{
    "1005/1005.4692_arXiv.txt": {
        "abstract": "We use \\textit{Chandra} X-ray and \\textit{Spitzer} infrared observations to explore the AGN and starburst populations of XMMXCS J2215.9-1738 at $z=1.46$, one of the most distant spectroscopically confirmed galaxy clusters known. The high resolution X-ray imaging reveals that the cluster emission is contaminated by point sources that were not resolved in \\textit{XMM-Newton} observations of the system, and have the effect of hardening the spectrum, leading to the previously reported temperature for this system being overestimated. From a joint spectroscopic analysis of the \\textit{Chandra} and \\textit{XMM-Newton} data, the cluster is found to have temperature $T=4.1_{-0.9}^{+0.6}$~keV and luminosity $L_{\\rm X} = (2.92_{-0.35}^{+0.24}) \\times 10^{44}$~erg~s$^{-1}$, extrapolated to a radius of 2~Mpc. As a result of this revised analysis, the cluster is found to lie on the $\\sigma_v-T$ relation, but the cluster remains less luminous than would be expected from self-similar evolution of the local $L_{\\rm X}-T$ relation. Two of the newly discovered X-ray AGN are cluster members, while a third object, which is also a prominent 24~$\\micron$ source, is found to have properties consistent with it being a high redshift, highly obscured object in the background. We find a total of eight $>5\\sigma$ 24~$\\micron$ sources associated with cluster members (four spectroscopically confirmed, and four selected using photometric redshifts), and one additional 24~$\\micron$ source with two possible optical/near-infrared counterparts that may be associated with the cluster. Examining the IRAC colors of these sources, we find one object is likely to be an AGN. Assuming that the other 24~$\\micron$ sources are powered by star formation, their infrared luminosities imply star formation rates $\\sim100$~M$_\\sun$~yr$^{-1}$. We find that three of these sources are located at projected distances of $< 250$~kpc from the cluster center, suggesting that a large amount of star formation may be taking place in the cluster core, in contrast to clusters at low redshift. ",
        "introduction": "\\label{s_intro} The universe at high redshift is a much more active place than we see locally. Radio and far-infrared studies have shown that the comoving star formation rate density increases by a factor $\\approx 10$ from $z=0$ to $z \\approx 1$ \\citep[e.g.][]{Seymour_2008, Magnelli_2009}, and continues rising towards higher redshifts, peaking at $z \\approx 3$ \\citep{Bouwens_2009}. The fraction of galaxies with Active Galactic Nuclei (AGN) is also seen to be much higher when the universe was young: current observations suggest that the space density of the brightest AGNs peaks at $z \\approx 2-3$ \\citep{Assef_2010, Silverman_2005}. Observations in the local universe show that the stellar mass of galactic bulges is correlated with the mass of nuclear super-massive black holes \\citep[e.g.][]{Ferrarese_2000, Gebhardt_2000}, suggesting that AGN play a major role in the growth of stellar mass. Simulations suggest that interactions and mergers between gas-rich galaxies drive this process, by providing both fuel for AGNs and triggering starburst activity in galactic nuclei \\citep[e.g.][]{Granato_2004, Hopkins_2008}. Eventually, gas heating by the AGN (feedback) reaches such a level that star formation is shut down, leaving the stellar population of the host galaxy to evolve passively. The quenching of star formation in massive galaxies by AGN feedback is a crucial ingredient in the latest semi-analytic models of galaxy formation, bringing the models into much closer agreement with observations \\citep[e.g.][]{DeLucia_2006, Croton_2006, Bower_2006, Somerville_2008}.  A connection between local environment, star formation and AGN is therefore expected in this theoretical picture. Intriguingly, recent studies of the dependence of star formation rate (SFR) on local galaxy density using large samples of field galaxies indicate that SFR increases with increasing local density at $z \\approx 1$ \\citep{Elbaz_2007, Cooper_2008} -- whereas the opposite is observed in the local universe \\citep[e.g.][]{Lewis_2002, Balogh_2004}. This situation is mirrored to some extent in galaxy clusters. Although these systems are dominated by quiescent, early type galaxies -- primarily `red-and-dead' ellipticals -- and continue to host substantial populations of these objects up to the highest redshifts observed \\citep[e.g.][]{Blakeslee_2003, Lidman_2008, Mei_2009, Hilton_2009}, an increasing fraction of star forming, late type galaxies are seen in clusters at higher redshifts \\citep[e.g][]{ButcherOemler_1984, vanDokkum_2000, Ellingson_2001, Smith_2005}. However, recent mid-infrared observations of clusters conducted using the \\textit{Spitzer Space Telescope} show that a significant amount of this increased star formation at high redshift is obscured by dust \\citep{Geach_2006, Marcillac_2007, Saintonge_2008, Haines_2009}, and therefore the amount of star formation in clusters has previously been underestimated.  The AGN fraction in clusters is also seen to rise with redshift. \\citet*{Martini_2009} report an eight-fold increase in the number of AGNs in clusters at $z=1$ compared to in the local universe. Similarly, \\citet{Galametz_2009} find that IR-selected clusters in the IRAC Shallow Cluster Survey \\citep{Eisenhardt_2008} show a clear excess of AGNs at high redshift, with the AGN fraction increasing by a factor of three from $0.5 < z < 1.0$ to $1.0 < z < 1.5$. However, it is not yet clear if the AGN fraction is higher in clusters than in the field, nor if the AGN fraction is evolving more rapidly in dense environments.  \\defcitealias{Stanford_2006}{S06} \\defcitealias{Hilton_2007}{H07} \\defcitealias{Hilton_2009}{H09}  In this paper we use \\textit{Chandra} X-ray observations and mid-infrared observations taken using the \\textit{Spitzer Space Telescope} to characterize the AGN and dusty star forming population of the $z=1.46$ cluster XMMXCS J2215.9-1738 (hereafter J2215), which was discovered in the ongoing optical follow-up of the \\textit{XMM} Cluster Survey \\citep[XCS;][]{Romer_2001, Sahlen_2009}. J2215 is the second highest redshift X-ray selected cluster known, following the recent discovery of a $z=1.62$ cluster in the Subaru/\\textit{XMM-Newton} Deep Field (\\citealt*{Tanaka_2010}, \\citealt{Papovich_2010}). The discovery of J2215 and an analysis of its X-ray properties using \\textit{XMM-Newton} data was reported in \\citet[][hereafter S06]{Stanford_2006}. We presented a first study of the dynamical state of the cluster and measurement of its velocity dispersion using additional VLT and Keck spectroscopy in \\citet[][hereafter H07]{Hilton_2007}. Most recently, we have used \\textit{Hubble Space Telescope} (\\textit{HST}) imaging and ground-based near-infrared data obtained at the \\textit{Subaru} telescope to perform a detailed examination of the morphologies of the cluster galaxies and the red-sequence within $<0.75$~Mpc of the cluster core \\citep[][hereafter H09]{Hilton_2009}.  The structure of this paper is as follows. In Section~\\ref{s_XRayAnalysis}, we report revised measurements of the cluster X-ray properties obtained through a joint spectroscopic analysis of the \\textit{Chandra} and \\textit{XMM-Newton} data, taking into account the effect of X-ray point sources that were unresolved in the \\textit{XMM-Newton} data used in \\citetalias{Stanford_2006}. In Section~\\ref{s_imaging} we briefly review the optical and ground based near-infrared data used in this paper, before describing new photometry obtained using the Infrared Array Camera \\citep[IRAC;][]{Fazio_2004} that extends wavelength coverage of the cluster into the rest frame near-infrared. We describe additional spectroscopic observations obtained at the Keck and Gemini observatories during 2008-2009 in Section~\\ref{s_OpticalSpectroscopy}, where we also present an updated measurement of the cluster velocity dispersion. Mid-infrared (24~$\\micron$) observations of the cluster obtained using the Multiband Imaging Photometer for \\textit{Spitzer} (MIPS) are described in Section~\\ref{s_MIPSAnalysis}. We characterize the properties of cluster 24~$\\micron$ and X-ray sources in terms of star formation or AGN activity in Section~\\ref{s_starFormation}. Finally, we discuss our findings in comparison with lower redshift studies of clusters in Section~\\ref{s_discussion}.  We assume a concordance cosmology of $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, and $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$ throughout, where $\\Omega_\\Lambda$ is the energy density associated with a cosmological constant. All magnitudes are on the AB system \\citep{Oke_1974}, unless otherwise stated. ",
        "conclusions": "\\label{s_conclusions} We have explored the AGN and star forming populations of the cluster J2215.9-1738 at $z=1.46$ using high resolution X-ray data from the \\textit{Chandra} satellite and infrared observations from the \\textit{Spitzer Space Telescope}. This is the first study of star formation as traced by mid-infrared observations in a cluster at $z \\approx 1.5$. We found:  \\begin{enumerate}  \\item{The cluster emission is contaminated by X-ray point sources, leading to the X-ray temperature being overestimated in the analysis presented in \\citetalias{Stanford_2006}; however, these point sources only contribute $\\approx 15$\\% to the total flux. Two of the X-ray point sources revealed by the \\textit{Chandra} observations are cluster members, while a third has properties consistent with it being a high redshift, highly obscured AGN in the background. The cluster has temperature $T=4.1_{-0.9}^{+0.6}$~keV and bolometric luminosity $L_{\\rm X} = (2.92_{-0.35}^{+0.24}) \\times 10^{44}$~erg~s$^{-1}$ (extrapolated to 2~ Mpc radius) from a joint analysis of the \\textit{XMM-Newton} and \\textit{Chandra} data. The cluster is less luminous than expected from self similar evolution of the local $L_{\\rm X}-T$ relation at the $\\approx 2\\sigma$ level.}  \\item{The velocity dispersion of the cluster is measured to be $\\sigma_v = 720 \\pm 110$~km~s$^{-1}$, from 31 galaxies within $R_{200}$. There is no clear evidence that the cluster velocity distribution is composed of two kinematically distinct components, although the \\citet{Hartigan2_1985} dip test of unimodality still hints that the cluster is not completely relaxed. Following the revised X-ray analysis presented in this paper, the cluster is found to lie on the $\\sigma_v-T$ relation, contrary to the result reported in \\citetalias{Hilton_2007}.}  \\item{Mid-infrared imaging reveals a total of eight $>5 \\sigma$ 24~$\\micron$ sources that are cluster members, selected using spectroscopic or photometric redshifts. In addition, there is a prominent 24~$\\micron$ source (J221559.68-173758.9) with two possible optical/near-infrared counterparts located $\\approx 17\\arcsec$ from the cluster center that may be associated with the cluster. One of the 24~$\\micron$ sources is found to have the infrared colors expected of an AGN, and has an elliptical morphology. The remaining objects are most likely powered by star formation, and if this is the case have SFRs $\\sim 100$~M$_\\sun$~yr$^{-1}$, adopting the \\citet{CharyElbaz_2001} spectral templates and assuming that the \\citet{Kennicutt_1998} law holds at this redshift.}  \\item{The cluster member AGN identified in the X-ray and infrared observations are located within $\\lesssim 2 \\sigma$ of the cluster red-sequence, as are three (possibly four) 24~$\\micron$ sources assumed to be powered by star formation. Three of the 24~$\\micron$ sources are also found within a projected distance of $< 250$~kpc from the cluster center, suggesting that the core of J2215 may be host to galaxies with very high star formation rates, in contrast to clusters at lower redshift.}  \\end{enumerate}"
    },
    "1005/1005.1886_arXiv.txt": {
        "abstract": "{ Over the past decade, astronomers have been using an increasingly\\ larger number of web-based applications and archives to conduct their\\ research.  However, despite the early success in creating links across\\ projects and data centers, the promise of a single integrated digital\\ library environment supporting e-science in astronomy has proven\\ elusive.  While some of the issues hampering progress in this area are\\ of technical nature, others are rooted in existing policies which\\ should be re-analyzed if further rapid progress\\ is to be made in this area. This paper describes a proposal that the\\ NASA Astrophysics Data System project has put forth in order to\\ improve its role as one of the primary discovery portals for astronomers,\\ focusing on those aspects which could benefit from an increased level\\ of involvement from the community, namely the effort to expose\\ astronomy resources as linked data, and the harvesting of\\ observational metadata.\\ }  \\FullConference{Accelerating the Rate of Astronomical Discovery, sps5\\\\ August 11-14, 2009\\\\ Rio de Janeiro, Brazil}  \\begin{document} ",
        "introduction": "Astronomy is an observational science. Our theories of the behavior of the universe and the objects in it are inspired by observations of that behavior, and are supported (or rejected) by observations, as well. Within the past twenty years systems of dense, interlinked data have become common and expected, fundamentally changing the way people think and services operate. It is now expected, for example, that from a listing of a movie at our local theater one can find a synopsis, reviews, a cast list, a list of all the director's films, recommendations, and the option to purchase tickets, as well.  Astronomy was one of the first disciplines to benefit from the early developments of web-based technologies enabling cross-linking of resources across archives (Accomazzi et al, 1994). Sixteen years ago, thanks to a collaboration between the NASA Astrophysics Data System (ADS) and the Centre de Donn{\\'e}es Astronomiques de Strasbourg (CDS), it became possible to go from a list of articles to the abstract of an article to a list of astronomical objects described in that article to a set of measurements on one of those objects. Thirteen years ago, again thanks to a collaboration between the ADS and several major data centers, including NED, HEASARC, MAST, ESO, and Chandra, it became possible to go from an article abstract to the actual observational data used to write the article, and then back to all publications describing each observation.  These connections have enabled astronomers to use the search capabilities of any of the main archives to locate a dataset or publication of interest, and then follow the appropriate links to find related information provided by another archive.  For instance, using ADS, one could search for papers on X-rays emission in Abell clusters and filter results to obtain just papers which have links to data products. While this is a very useful way to narrow bibliographic searches, selecting which of these papers has links to images or spectra or catalogs still requires a person to click through all the data links provided in the list of returned papers. Automating this activity is currently not practical since connections between one archive and another are purely defined as links between URLs, and lack semantic and contextual information between the resources they represent.  In this paper we argue that a tighter integration of observational and bibliographic metadata will enable the creation of new connections between resources in astronomy and applications that will allow users to search, browse, and reason over them.  The existence of these connections between and within papers and data products represents more that just a convenient feature to the end-user.  The research process in science today involves the generation, retrieval, manipulation, analysis, and publication of digital artifacts which are typically stored in different archives on the web.  In order to enable scientists to access these scholarly products, and to document and recreate the analysis that was carried out on them, a formal description of these resources and processes is necessary (Pepe et al., 2009). Establishing relationships and links among these resources provides a way to satisfy a number of different goals, among them: aggregation (linking together all artifacts used in a study, including data products, notes, draft papers, software tools); attribution (properly acknowledging prior work, tools being used, or datasets being analyzed); preservation (maintaining provenance information and reconstructing the workflow used in the research process); and discovery (following connections between resources may uncover previously unknown relationships). ",
        "conclusions": "In this coming era of data intensive science, it is increasingly important to be able to seamlessly move between scientific results, the data used to publish them, and the processes used to produce them.  In addition, scientific research requires that we are able to establish the provenance of data sources and processes operating on this data, so that research may be repeated, variations on the research be carried out, and new research on existing datasets be enabled.  While the Virtual Observatory has provided us with the standards and protocols needed to model and exchange astronomical datasets, the problem of information discovery in our data-centric world still looms large.  The key to creating such a system is having access to the metadata describing the datasets at the appropriate granularity level so that the proper resources can be identified, described, and connections between them can be made.  The technology which can be used to expose and link these metadata has now matured and is the foundation of the semantic web framework and in particular in the linked data model. These technologies promise to become widely-adopted standards enabling the creation and growth of a ``web of data'' in which machine-readable content parallels the human-readable content available today.  We believe that the astronomical community would benefit from embracing this philosophy and adopting these technologies. In particular, three sets of metadata resources currently stored in our archives should be made available machine-readable format for harvesting and indexing: observing proposals, their attributes and links to observations; observations and their attributes; instruments, their attributes and capabilities. From a technical standpoint this metadata should be exposed as RDF and should make use of community-developed (lightweight) ontologies.  This will allow the use of linked-data principles to properly connect these resources. However, even before these technical aspects are worked out, we believe that the availability of the metadata to organizations such as the ADS, the CDS, NED, and other IVOA members will enable the creation of knowledge bases and applications demonstrating the value of such an interlinked data system. Therefore we urge the community to support this effort by recommending that observatories and archives make available and permit the harvesting and indexing of metadata describing publicly accessible datasets, including observing proposals, observing logs, and FITS headers."
    },
    "1005/1005.3551_arXiv.txt": {
        "abstract": "{ Massless interacting scalar fields in de Sitter space have long been known to experience large fluctuations over length scales larger than Hubble distances. A similar situation arises in condensed matter physics in the vicinity of a critical point, and in this better-understood situation these large fluctuations indicate the failure in this regime of mean-field methods. We argue that for non-Goldstone scalars in de Sitter space, these fluctuations can also be interpreted as signaling the complete breakdown of the semi-classical methods widely used throughout cosmology. By power-counting the infrared properties of Feynman graphs in de Sitter space we find that for a massive scalar interacting through a $\\lambda \\, \\phi^4$ interaction, control over the loop approximation is lost for masses smaller than $m \\simeq \\sqrt \\lambda \\, H/2\\pi$, where $H$ is the Hubble scale. We briefly discuss some potential implications for inflationary cosmology.}  \\begin{document} ",
        "introduction": "Some aspects of quantum fields on de Sitter space remain controversial, long after their first investigation more than 30 years ago \\cite{dSorig}, but their potential relevance as an explanation of the detailed properties of the fluctuations observed by precision measurements \\cite{CMBmeas, CMBth1} in the Cosmic Microwave Background (CMB) radiation has stimulated a recent re-examination of the issues \\cite{Wbg, dSrecent1, dSrecent2, KLR,AKH, Dolgov, MS, petri, us, SZ} (see also \\cite{Review} and references therein)  The main interpretational difficulties for de Sitter space arise for massless fields, or for those that are very light compared with the Hubble scale: $m \\ll H$. There are two related ingredients that complicate calculations with such fields: the presence of various types of infrared singularities, and the presence of large fluctuations over extra-Hubble distances.  For example, if a massless scalar field is prepared in an initial state for which fluctuations are small, then quantities like $\\langle \\phi^2(t) \\rangle$ grow linearly with cosmic time, $t$. For massive scalars this growth eventually saturates at a $t$-independent value $\\langle \\phi^2 \\rangle \\propto H^4/m^2$, which is parametrically large if $m \\ll H$ \\cite{lint}. Since these fluctuations are uncorrelated, $\\langle \\phi(x) \\phi(y) \\rangle \\simeq \\langle \\phi(x) \\rangle \\langle \\phi(y) \\rangle$, on scales longer than $H^{-1}$ and since gradients quickly redshift away, one finds a picture in which the field takes an approximately constant value within each Hubble volume, with different volumes evolving independently of one another in an essentially random and uncorrelated way \\cite{superhubble}.  Although this has long been recognized as the appropriate physical picture, essentially all of what we know about fields in de Sitter backgrounds is based on calculations performed within the semi-classical approximation. This approximation assumes that classical field theory captures the dominant physics, with small calculable corrections arising from quantum fluctuations. The size of these corrections is believed to be kept small because of their systematic dependence on small dimensionless quantities, like powers of coupling constants, $\\lambda/(4\\pi)^2$, and (because gravity is non-renormalizable) of small energy ratios, $E/M_p$ \\cite{GREFT}. In particular, such an approximation underpins calculations of inflaton fluctuations during inflation, and their implications for the properties of the CMB at recombination.  In this paper we argue that for massless scalar fields in de Sitter backgrounds subject to non-derivative self-interactions --- like $V_{\\rm int} \\simeq \\lambda \\, \\phi^4$ --- the presence of large extra-Hubble fluctuations undermines the entire semi-classical approximation. This is because the semiclassical approximation is at heart a mean-field approximation, within which a quantum field is represented as a dominant classical configuration plus a small quantum fluctuation, $\\phi(x) = \\varphi(x) + \\delta\\phi(x)$. But this kind of a description fundamentally breaks down over distances larger than $H^{-1}$ due to the large fluctuations occurring on these scales.  Notice that the breakdown of semiclassical methods we have in mind does not merely mean that the classical approximation is inadequate, with the situation being saved if we compute just a few more loops than usual. Instead, for massless scalars the danger is that the assumption that higher loops are suppressed by a small quantity that breaks down, meaning that {\\em all} orders in the semiclassical expansion have similar sizes. This makes semiclassical calculations inherently unreliable because the truncation of the loop expansion omits contributions that are as large as those that are kept.  Quantum field theory is a cruel but fair master, so (as usual) the formalism contains within itself the news about the breakdown of semiclassical methods. The messenger is in this case the infrared divergences that commonly plague de Sitter calculations with massless scalars. These indicate a singular dependence on the mass in the more general case of a massive, but very light, scalar field. We argue that at least some of these divergences reflect the dominance of fluctuations over the contributions of the classical background, pointing to a fundamental breakdown of mean-field methods.  To make our claim precise, we consider a scalar field in de Sitter space, that self-interacts through a quartic scalar potential: $V = \\frac12 \\, m_0^2 \\phi^2 + \\frac{1}{4!} \\, \\lambda \\, \\phi^4$. Perturbing in $\\lambda$ and using propagators for fields of mass $m_0$ shows that in the small-mass limit the usual loop-suppressing factor of $[\\lambda/(2\\pi)^2]$ for each loop in an $L$-loop graph is systematically modified by factors of $(H^2/m_0^2)$, indicating that higher loops are not suppressed once the scalar mass is sufficiently light.  We explicitly identify contributions to $L$-loop correlators that are proportional to \\be \\left( \\frac{\\lambda H^4}{4 \\pi^2 m_0^4} \\right)^{L} \\,, \\ee which, taken at face value, would indicate that perturbation theory fails once $m_0^2 \\sim \\sqrt\\lambda \\, (H/2\\pi)^2$. However we argue that the breakdown of perturbative methods at this higher mass arises because the physical mass scale that cuts off IR effects is really $m^2 = m_0^2 + \\delta m^2$, with $\\delta m^2 \\simeq \\lambda H^4/m_0^2$, and so it is only the expansion $H^2/m^2 = (H^2/m_0^2) [1 - \\delta m^2/m_0^2 + \\cdots]$ that breaks down when $m_0^2 \\sim \\sqrt\\lambda \\, (H/2\\pi)^2$. This particular breakdown can be resummed: that is, it can be removed by reorganizing the perturbative expansion so that the unperturbed lagrangian involves the mass term $m^2 \\phi^2$ rather than $m_0^2 \\phi^2$. This particular reorganization does not also remove the potential breakdown of perturbation theory at $m_0^2 \\simeq \\lambda (H/2\\pi)^2$.  Finite-temperature field theory provides a well-understood precedent for these conclusions. The small-momentum limit of the Bose-Einstein distribution function, $n_\\ssB(k) \\propto T/k$, implies that infrared divergences are stronger at finite temperature than they are at zero temperature. As a result, an $L$-loop Feynman graph for a self-interacting scalar field at finite temperature comes with the systematic factor $(\\lambda \\, T/4\\pi^2 m_0)^{L}$, again indicating a breakdown of the loop expansion if $m_0 \\lsim \\lambda \\,T/(2\\pi)^2$. In the thermal case it is also known that the loss of perturbative control is only partial when $m_0 = 0$ since then the complete thermal mass is $m^2 =  \\delta m^2 \\propto \\lambda T^2$ and the IR divergences can be weakened by reorganizing perturbation theory so that the unperturbed fields have this mass. In this case the same power-counting leads to the loop factor $\\lambda T/m \\propto \\sqrt \\lambda$, leading to controlled (but non-analytic) results for small $\\lambda$. This same resummation fails, however, if $m_0^2 < 0$ is adjusted so that $m^2 = 0$, leading to a {\\em bona fide} breakdown of semiclassical methods, such as is known to occur in condensed matter systems in the vicinity of a critical point. The breakdown of the loop expansion in this case underlies the well-known failure of mean-field methods to describe critical exponents \\cite{RGCrit}. The physical origin of this breakdown is the dominance of large fluctuations near the critical point, similar to the fluctuations found in de Sitter space.  The analogy between thermal field theory and de Sitter space is robust, apart from one potentially important difference: in de Sitter space the difference $\\delta m^2 = m^2 - m_0^2 \\propto \\lambda H^4/m_0^2$ itself depends singularly on $m_0^2$, unlike for thermal field theory where $\\delta m^2 \\propto \\lambda T^2$. Because of this difference the map between the two cases is not simply to think of the Hubble scale as a temperature. In particular, although $m^2$ can be adjusted to vanish in the thermal case by appropriately choosing $m_0^2$, it is not clear that this can be done in the de Sitter situation. We discuss in the conclusions the open problem  of the extent to which these same regimes of resummation also apply in the de Sitter case (however, see \\cite{Akhmedov:2008pu}).  Although gravity resembles a massless scalar in many ways, we emphasize that we do {\\em not} expect this same failure to arise for pure gravity in a de Sitter background. The difference arises because the gravitational self-interactions are largely derivative couplings, and so are typically not as divergent in the infrared. Since Goldstone scalars similarly couple derivatively, they also need not share the same mean-field breakdown. Massless spin-1 fields can couple without derivatives, and infrared effects are also known to ruin mean-field methods for a hot plasma of charged particles interacting through gauge interactions \\cite{HTL}. We leave it open whether a similar breakdown occurs on de Sitter space, but any such a failure would require the existence of very light charged degrees of freedom to survive the exponential de Sitter red-shifting. There is some work for SQED in de Sitter space that shows that the photon can indeed get a mass there (see \\cite{PW} for a review).  The remainder of the paper is organized as follows. The next section, \\S2, briefly reviews the power-counting of infrared divergences for self-interacting scalar fields at finite temperature, to show how these track the breakdown of mean-field methods. \\S3 then provides a similar power-counting for self-interacting scalars in de Sitter space. Dimensional analysis is first used to argue that the worst divergences are logarithmic for any correlation functions. For a massive propagator, this translates into inverse powers of mass at each loop order. A class of graphs is then displayed that verifies this dependence through explicit calculation. Finally, \\S4\\ draws some preliminary implications for inflationary calculations, and summarizes our conclusions. ",
        "conclusions": "The body of this paper argues that an $L$-loop contribution to a correlation function for a scalar field in de Sitter spacetime with $\\lambda \\, \\phi^4$ self-interactions carries a systematic factor of $(\\lambda H^2/4\\pi^2 m_0^2)^{L}$, indicating a fundamental breakdown of semiclassical methods once $m_0 \\lsim \\sqrt\\lambda \\; H/2\\pi$. The origin of the breakdown of perturbative methods is the infrared-singular behavior of these graphs which arises due to the large extra-Hubble fluctuations experienced by very light scalars in de Sitter space. These fluctuations dominate the contributions to correlation functions, invalidating the semiclassical approximation which is at its heart a mean-field description.  This story is qualitatively similar to what happens in finite temperature field theory, although the power of coupling constant that defines the boundary of the semiclassical region differs. In the finite temperature case, the variance of the field, $\\Expect{\\phi^2}$, goes like $T^ 2$ while in the de Sitter case we find $\\Expect{\\phi^2}\\sim H^4/m^2$. This difference means that the mass, $m_{\\rm dyn}$, which is comparable to the corrections to $m$ scales differently with $\\lambda$ in these two cases: for the thermal case $m_{\\rm dyn}^2 \\propto \\lambda T^2$, while for a de Sitter background $m_{\\rm dyn}^2 \\propto \\lambda^{1/2} H^2$. It also may mean that the structure of de Sitter space precludes access to the regime $m^2 \\simeq 0$ for any choice of $m_0^2$, unlike for finite-temperature systems, see Figure (\\ref{fig:range}). However, as is shown above, the two theories have a similar perturbative structure, whose relation is sketched in Figure (\\ref{fig:chart}).   \\FIGURE[ht]{ \\epsfig{file=massplot.eps,angle=0,width=0.5\\hsize} \\caption{A plot of $m^2$ vs $m_0^2$. The horizontal band represents the regime $m^2 < \\cO(\\lambda H^2/4\\pi^2)$ while the curve is given by Eq. (\\ref{massshift}). It would be interesting to investigate the case of negative bare mass, which we do not display here, further} \\label{fig:range}}  A natural question to ask about this perturbative breakdown at small $m^2$ is whether it can be resummed (like for $m_0^2 \\ge 0$ at finite temperature), or whether it reveals a complete breakdown of expansions based on powers of $\\lambda$ (like when $m_0^2 < 0$ is chosen so that $m^2 = 0$ at finite temperature).  \\FIGURE[ht]{ \\epsfig{file=MassRangeFigSmall.eps,angle=0,width=0.6\\hsize} \\caption{A comparison of the behavior of loop corrections for a scalar field with a quartic interaction $\\lambda\\, \\phi^4$ in a thermal background or de Sitter space as a function of mass. The masses labelled on the chart are, from left to right, the mass below which perturbation theory breaks down, the dynamical mass generated in the event that the zero temperature or flat space mass was zero, and the maximum mass the field can have and still receive thermal/de Sitter fluctuations.} \\label{fig:chart} }  Others have argued that there exist semiclassical methods that capture the leading infrared logs. For secular growing logs, one proposal solves the late-time physics using the classical equation of motion \\cite{MS}, in a similar way to the $\\delta N$ formalism. Another point of view uses a stochastic approach to inflation \\cite{superhubble}, which is argued to capture (and resum) the leading logs, by generating a dynamical mass \\cite{SY}. Although such a stochastic approach goes beyond mean field, it is not yet clear what combination of small parameters control the approximations made in its derivation.  {}From the point of view of the arguments made here, derivative couplings are not as dangerous as are those of the scalar potential. This is because the momentum dependence of these couplings tends to ameliorate any IR divergence the graph would have otherwise had. This means that massless Goldstone bosons on de Sitter space would {\\em not} suffer from the same breakdown of perturbation theory as does the $\\phi^4$ model considered here.  We would expect the self-interactions of massless gravitons to be similarly benign so long as these are derivative couplings, leading us to expect no perturbative breakdown for pure gravity on de Sitter space. This expectation seems to be borne out by ref.~\\cite{DB} and ref.~\\cite{GS} (the latter appeared just as our paper neared completion). These authors study one-loop infrared divergences for pure gravity and gravity coupled to scalars in de Sitter and slow-roll spacetimes, but find that IR divergences cancel in the absence of the self-interactions of the scalar potential (which appear in their calculations as slow-roll parameters). Based on the power-counting arguments presented here we expect this to continue to be true at higher loops, with the scalar self-interactions being the most dangerous in the infrared.  It is interesting to consider in this light what implications our result might have for inflationary cosmologies. A complication in directly extracting these for simple single-field models is our neglect of classical evolution of the background scalar field and metric, due to our use of a simple de Sitter background. Because a homogeneous evolving scalar field can be used to define a notion of cosmic time, many of its effects can be gauged away. We therefore expect a naive application of the above arguments to simple models are likely to cancel from gauge invariant quantities like the curvature fluctuations of physical interest for cosmology. Nevertheless, it may be possible to have important infrared effects appear in curvature correlations in multi-field models, particularly those involving nontrivial post-inflationary dynamics (such as curvaton models). We leave for future work the detailed question of whether and how the infrared effects we find above `propagate' into late-time curvature perturbations.  In the remainder of this section we put these issues aside, and ask what the condition $m^2 > \\lambda H^2/4\\pi^2$ implies for the parameters of a single scalar field described by a quartic potential \\be V = V_0 + \\frac12 \\, m^2 \\phi^2 + \\frac{1}{4!} \\, \\lambda\\, \\phi^4\\,. \\ee For this model we ask how the condition $M^2 \\gg \\lambda (H/2\\pi)^2$, where $M^2 = V'' = m^2 + \\frac12 \\lambda \\, \\phi^2$, compares to other conditions to which inflationary models must be subject. To see what this implies for potential inflationary applications, consider two extreme cases: $(i)$ $\\phi$ so large that $V \\simeq \\frac{1}{24} \\, \\lambda \\, \\phi^4$; and $(ii)$ $\\phi$ small enough that $V \\simeq V_0$. (This last example can be regarded as part of a model of hybrid inflation \\cite{hybrid}, with inflation ending as another field starts to roll as its $\\phi$-dependent squared-mass goes negative.)  \\subsubsection*{Large-field inflation}  For the large-$\\phi$ regime we have $V \\simeq \\frac{1}{24} \\, \\lambda \\, \\phi^4$, and so $H^2 = V/3M_p^2 \\simeq \\lambda \\, \\phi^4/72 M_p^2$. The slow-roll parameters are \\be \\varepsilon := \\frac12 \\left( \\frac{M_p V'}{V} \\right)^2 \\simeq 4 \\left( \\frac{M_p}{\\phi} \\right)^2 \\quad \\hbox{and} \\quad \\eta := \\frac{M_p^2 V''}{V} \\simeq 12 \\left( \\frac{M_p}{\\phi} \\right)^2  \\,, \\ee so the edge of the inflationary regime is $\\phi_\\SR /M_p \\simeq \\O(1)$. For $\\phi$ larger than this classical evolution satisfies \\be \\dot\\phi \\simeq \\frac{V'}{3H} \\simeq \\frac{\\frac16 \\, \\lambda \\, \\phi^3}{\\sqrt\\lambda \\; \\phi^2/2\\sqrt2 \\; M_p} \\simeq \\frac{\\sqrt{2}\\,}{3} \\, \\sqrt\\lambda \\; M_p \\phi \\,. \\ee Once $\\dot\\phi$ becomes smaller than $H^2$ fluctuations dominate classical evolution and inflation becomes eternal, which in this case occurs when $\\frac{\\sqrt{2}\\,}{3} \\sqrt\\lambda \\; M_p \\phi < \\lambda \\, \\phi^4/72 M_p^2$, or $\\phi^3 > \\phi_\\EI^3 \\simeq 24 \\sqrt2 \\; M_p^3/\\sqrt\\lambda$.  How do the boundaries of the semiclassical approximation compare to these? There are two criteria to be satisfied. First, control over the low-energy approximation that underlies the gravitational loop expansion requires $V \\ll M_p^4$, or $\\phi^4 \\ll \\phi_\\HE^4 \\simeq 24 M_p^4/\\lambda$. We have seen in previous sections that the $\\lambda$ loop expansion fails unless $M^2 \\gg \\lambda \\, H^2/4 \\pi^2$ or $\\frac12 \\, \\lambda \\, \\phi^2 \\gg \\lambda^2 \\phi^4/288 \\pi^2 M_p^2$, and so $\\phi^2 \\ll \\phi^2_\\BD \\simeq (12 \\pi M_p)^2/\\lambda$. Since $\\phi_\\SR/M_p \\simeq \\O(1)$, $\\phi_\\EI/M_p \\simeq \\O(\\lambda^{-1/6})$, $\\phi_\\HE /M_p \\simeq \\O( \\lambda^{-1/4})$ and $\\phi_\\BD /M_p \\simeq \\O( \\lambda^{-1/2})$ we have \\be \\phi_\\SR \\ll \\phi_\\EI \\ll \\phi_\\HE \\ll \\phi_\\BD \\,, \\ee and so the condition $M^2 > \\lambda (H/2\\pi)^2$ is parametrically weaker than the condition $M_p^4 > V$. The fact that eternal inflation can occur before a total breakdown of perturbation theory is as expected.  \\subsubsection*{Hybrid inflation}  Consider next a hybrid model \\cite{hybrid}, involving two scalar fields, $\\phi$ and $\\chi$, interacting through the potential \\be U(\\phi,\\chi) = \\frac{1}{4} \\,\\zeta \\,( \\chi^2 - v^2 )^2 + \\frac{g^2}{2} \\, \\chi^2 \\phi^2 + \\frac12 \\, m^2 \\phi^2 + \\frac{1}{4!} \\, \\lambda\\, \\phi^4\\;. \\ee In this model the fields start in the trough defined by $\\chi = 0$, with $\\phi$ large and rolling towards smaller values subject to the effective potential \\be V = V_0 + \\frac12 \\, m^2 \\phi^2 + \\frac{1}{4!} \\, \\lambda\\, \\phi^4\\,, \\ee with $V_0 = \\frac{1}{4!} \\, \\zeta \\, v^4$. This roll continues until \\be \\phi = \\phi_\\SR = \\frac{\\sqrt\\zeta \\, v}{g} \\,, \\ee at which point the $\\chi$ mass, $\\mu^2 = - \\zeta \\, v^2 + g^2 \\phi^2$, becomes negative, allowing $\\chi$ to evolve quickly towards the absolute minimum at $\\chi = v$ and $\\phi = 0$.  The dynamics of the inflaton in this model is governed by the same potential considered earlier, but our interest now is in the small-field regime, for which $m^2 \\gg \\frac12 \\, \\lambda \\, \\phi^2$ and $V \\simeq V_0$ (and so $H^2 \\simeq V_0/3M_p^2$). We assume parameters are chosen to keep $\\phi_\\SR$ small enough to ensure $\\chi$ remains zero well into this regime. In this case the slow-roll parameters are \\be \\varepsilon \\simeq \\frac12 \\left( \\frac{m^2 M_p \\phi}{V_0} \\right)^2  \\quad \\hbox{and} \\quad \\eta \\simeq \\frac{m^2 M_p^2}{V_0}  \\,, \\ee so $2 \\varepsilon \\simeq \\eta^2 (\\phi/M_p)^2$. Clearly inflation only requires $m^2 M_p^2 /V_0 \\ll 1$, since the condition $\\frac12 \\, \\lambda \\phi^2 \\ll m^2$ automatically ensures $\\phi \\ll M_p/\\eta$. Inflation ends once $\\phi$ reaches $\\phi_\\SR = \\sqrt\\zeta \\; v/g$.  In this case the validity of the $\\lambda$ loop expansion requires $m^2 \\gg \\lambda \\, H^2/4 \\pi^2 \\simeq \\lambda V_0/12 \\pi^2 M_p^2$, or \\be \\eta \\simeq \\frac{m^2 M_p^2}{V_0} \\gg \\frac{\\lambda}{12 \\pi^2} \\,. \\ee Notice that this lower bound to the slow-roll parameters is not simply a naturality condition, corresponding to a regime for which small loop corrections dominate the smaller classical potential. Rather, in this regime the problem is not fixed simply by including one- or two-loop corrections to $V$. Instead it is the entire loop expansion itself that fails. It would be of great interest to see whether the semiclassical criterion plays any role in more general inflationary contexts."
    },
    "1005/1005.2411_arXiv.txt": {
        "abstract": "{We have determined the luminosity function of 250$\\,\\mu$m-selected galaxies detected in the $\\sim$14\\,deg$^2$ science demonstration region of the Herschel-ATLAS project out to a redshift of $z=0.5$. Our findings very clearly show that the luminosity function evolves steadily out to this redshift. By selecting a sub-group of sources within a fixed luminosity interval where incompleteness effects are minimal, we have measured a smooth increase in the comoving 250$\\,\\mu$m luminosity density out to $z=0.2$ where it is $3.6^{+1.4}_{-0.9}$ times higher than the local value.} ",
        "introduction": "Measurement of the galaxy luminosity function (LF) constitutes one of the most fundamental statistical constraints that can be placed on models of galaxy formation, and hence the build up of large scale structure, in the universe. Since half of the energy ever emitted by galaxies has been absorbed by dust and re-radiated at far-infrared and sub-millimetre (submm) wavelengths (Fixsen et al. \\cite{fixsen98}) and, because knowledge of the statistical properties of submm sources is relatively sparse, determination of the submm LF provides a crucial missing piece in a fully comprehensive model of galaxy evolution.  Following their detection in the first deep submm and mm surveys (e.g., Smail et al. \\cite{smail97}; Hughes et al. \\cite{hughes98}; Eales et al. \\cite{eales99}; Bertoldi et al. \\cite{bertoldi00}), much has been learned about the dusty high-redshift sources selected at such wavelengths. Although many studies have argued that these sources are likely ancestors of local ellipticals (e.g., Scott et al. \\cite{scott02}; Dunne et al. \\cite{dunne03}) little progress in verifying this assertion has been made since.  The main reason for this is the preponderance of high redshift submm/mm-selected sources, owing to the strong negative k-correction and small survey areas. The resulting low numbers of sources at redshifts $z<1$ has therefore precluded evolutionary studies over the last $\\sim$60\\% of the Universe's history.  In particular, despite the local submm LF being first determined a decade ago (Dunne et al. \\cite{dunne00}), little has been added to our comprehension of how the LF has evolved over the last $\\sim 7$\\,Gyr, until very recently. Observations conducted using the Balloon-borne Large Aperture Submm Telescope (BLAST; Devlin et al. \\cite{devlin09}) have made significant improvements with much enhanced sensitivity to $z<1$ sources. As a result, direct estimates of the LF at 250, 350 and 500$\\,\\mu$m were made by Eales et al.  (\\cite{eales09}) who detected strong evolution, particularly among the higher luminosity systems, from $z=1$ to the present day. However, the accuracy of these findings is limited by the small number of sources used ($\\sim 50$ at $z<0.5$) and source confusion due to the angular resolution of BLAST.  In this letter, we present our measurement of the LF of 250$\\,\\mu$m-selected galaxies detected by the Herschel Space Observatory ({\\sl Herschel}; Pilbratt et al. \\cite{pilbratt10}) over a $\\sim 14\\,$deg$^2$ region acquired as part of the science demonstration observations of the {\\sl Herschel}-Astrophysical Terahertz Large Area Survey (H-ATLAS; Eales et al. \\cite{eales10}). These data offer a significant improvement over the BLAST data in terms of their increased sensitivity, higher angular resolution and greater areal coverage, resulting in $\\sim 20$ times the number of sources with which to compute the LF.  Throughout this letter, the following cosmological parameters have been assumed; ${\\rm H}_0=71\\,{\\rm km\\,s}^{-1}\\,{\\rm Mpc}^{-1}$, $\\Omega_{\\rm m}=0.27$, $\\Omega_{\\Lambda}=0.73$. ",
        "conclusions": "One of the key goals of H-ATLAS will be to understand the nature of the evolution detected in this letter. In turn, we aim to improve our understanding of the evolutionary link between high redshift and local submm systems.  This letter has only considered sources selected at 250$\\,\\mu$m, merely one of the five wavebands on offer from H-ATLAS. Furthermore, the 14\\,deg$^2$ of survey data analysed in this work represent only 2.5\\% of the final, proposed H-ATLAS survey area. A repeat of the analysis presented here with the final survey data, would therefore result in the quoted uncertainties falling by at least a factor of five. In light of these considerations, it is clear that our results offer only a small glimpse of the anticipated wealth of science that H-ATLAS has to offer."
    },
    "1005/1005.0622_arXiv.txt": {
        "abstract": "In this article, we test the hypothesis that Cepheids have infrared excesses due to mass loss.  We fit a model using the mass-loss rate and the stellar radius as free parameters to optical observations from the OGLE-III survey and infrared observations from the 2MASS and SAGE data sets.  The sample of Cepheids have predicted minimum mass-loss rates ranging from zero to $10^{-8}M_\\odot$ $yr^{-1}$, where the rates depend on the chosen dust properties.  We use the predicted radii to compute the Period-Radius relation for LMC Cepheids, and to estimate the uncertainty caused by the presence of infrared excess for determining angular diameters with the infrared surface brightness technique.  Finally, we calculate the linear and non-linear Period-Luminosity (P-L) relations for the LMC Cepheids at $VIJHK$ + IRAC  wavelengths and we find that the P-L relations are consistent with being non-linear at infrared wavelengths, contrary to previous results. ",
        "introduction": "Classical Cepheids are powerful standard candles because they follow the  Period-Luminosity (P-L)  relation or Leavitt Law \\citep{Leavitt1908}. The P-L relation has been used as a tool for Galactic \\citep{Feast1997b}, extragalactic \\citep[and references therein]{Gieren2009} and cosmological studies \\citep{Freedman2001, Sandage1998}.  The LMC Cepheids have been central for deriving these relations, initially using optical data \\citep{Udalski1999a,Sandage2004,Kanbur2003, Kanbur2006, Fouque2007}, then near-infrared observations \\citep{Groenewegen2000, Persson2004, Fouque2007}, and recently using infrared observations \\citep{Freedman2008, Ngeow2008a}.  The infrared P-L relations were determined by matching sources from the SAGE (Surveying the Agents of a Galaxy's Evolution) archival data \\citep{Meixner2006} with the known Cepheids.  \\cite{Ngeow2008b} computed the infrared P-L relations for a larger sample of 1848 Cepheids from the OGLE-III Database \\citep{Soszynski2008} by matching those sources with the two published epochs of SAGE data, and also averaging the two epochs of infrared fluxes to bring the fluxes closer to the mean flux. They tested the P-L relations for non-linearity, and found that the IR P-L relations are linear, while the relations at wavelengths shorter than $K$-band are non-linear.  However, there are a number of Cepheids with infrared fluxes that are $>3$ standard deviations brighter than predicted by the P-L relations. Similar results were found using AKARI N-band data \\citep{Ngeow2010}. It is important to understand the nature of this excess infrared brightness because the intrinsic infrared flux of Cepheids is less metallicity dependent than optical P-L relations \\citep{Freedman2008}, which makes the IR P-L relations powerful tools for extragalactic studies.  The James Webb Space Telescope will be able to observe Cepheids in distant galaxies, making it possible to use the IR P-L relations to determine distances, but the unknown source of infrared excess will increase the uncertainty of the results.  One hypothesis for the source of the infrared excess is the existence of a circumstellar envelope (CSE) of dust around the LMC Cepheids, analogous to the CSEs observed around Galactic Cepheids \\citep{Kervella2006, Merand2006, Merand2007}.  Multi-wavelength observations of RS Puppis and $l$ Car confirm the presence of CSEs around these Cepheids \\citep{Kervella2009}. There is more evidence of infrared excess in Galactic Cepheids using IRAS observations \\citep{McAlary1986, Deasy1988} and Spitzer observations \\citep{Marengo2009}. On the other hand, \\cite{Marengo2009b} argued that there is no evidence from IR excess from warm dust based on Spitzer observations and instead was due to carbon monoxide emission.  The Cepheid RS Pup is also associated with a nebula \\citep{Havlen1972, Kervella2008}.  \\cite{Kervella2006} speculated that CSEs surrounding Galactic Cepheids were caused by mass loss that can leads to an infrared excess. \\cite{Neilson2008a} proposed  a driving mechanism for this wind that is a combination of radiative acceleration, pulsation and shocks in the atmosphere of the Cepheid. This proposed driving mechanism  was tested by developing an analytic mass-loss model that was used to calculate rates of mass loss for a sample of Galactic Cepheids, with values between $10^{-10}$ to $10^{-7}M_\\odot/yr$ being found.  This model predicts the fluxes of CSEs around the Cepheids that agree well with those observed.   \\cite{Neilson2008b} applied this mass-loss model to theoretical models of Galactic, LMC, and Small Magellanic Cloud Cepheids and found that Cepheids in all three galaxies have significant mass-loss rates and infrared excesses.   There is other evidence for Cepheid mass loss in addition to infrared excess; \\cite{Nardetto2008b} observed H$\\alpha$ line profiles of Galactic Cepheids with the purpose of detecting a hydrogen circumstellar medium.  They found that the H$\\alpha$ line profiles are asymmetric and blue-shifted in long period ($P>10$ day) Cepheids, possibly caused by stellar wind.  While it is important to characterize infrared excess of Cepheids in order to use them for various applications, it is just as important to understand how mass loss affects the structure and evolution of Cepheids.  Mass loss is a potential solution to the mass discrepancy of Cepheids, which is the difference between Cepheid masses estimated  using stellar evolution calculations and estimates using stellar pulsation calculations \\citep{Cox1980}.  Currently, stellar pulsation calculations predict masses about $10$-$20\\%$ smaller than stellar evolution models \\citep{Keller2006, Keller2008}. Measurements of dynamic masses of Cepheids in binary systems tend to agree with pulsation calculations \\citep{Evans2006, Evans2008}, suggesting the discrepancy is in the stellar evolution modeling.  If Cepheids lose an average of $10^{-8}$ to $10^{-7}M_\\odot/yr$, then mass loss could solve the mass discrepancy.  Mass loss is hypothesized to create optically thin circumstellar dust shells that cause infrared excess in Cepheids but do not affect the visual extinction of the Cepheids.  This excess should be seen in infrared observations of the LMC from the SAGE survey \\citep{Meixner2006} for which there are currently two epochs of observations.  In fact, infrared excess has been detected in a number of evolved AGB stars in the LMC \\citep{Vijh2008, Groenewegen2009}.    In our earlier paper testing mass loss in LMC Cepheids, we found that mass loss may be important, with about $10\\%$ of OGLE-II Cepheids showing statistically significant evidence for dust shells. However,  the uncertainty of mean infrared flux of the Cepheids is a limiting factor.  The predicted gas mass-loss rates of the LMC Cepheids in that sample range from $10^{-11}$ to $10^{-8}M_\\odot/yr$ assuming a dust-to-gas ratio of $1/250$.  Mass-loss rates in this range affect both the zero-point and the slope of the infrared P-L relations.  When the infrared excess was removed from the OGLE-II sample, \\cite{Neilson2008c} found that the scatter of the data was reduced. We also found that the $3.6$ to $5.8$ $ \\mu m$ relations are consistent with being non-linear while the $8.0$ $\\mu m$ relation are linear, differing from the observed relations \\citep{Ngeow2008a}.  The infrared excess caused by mass loss appeared to have a greater effect at shorter periods than at longer periods.  However, the infrared stellar luminosity is an increasing function of period, making the infrared excess is less noticeable for long-period Cepheids, even for larger mass-loss rates  ($\\approx 10^{-8}~M_\\odot/yr$).  These results suggest that the near-infrared fluxes, $JHK$, are potentially affected by mass loss.  This would mean distance estimates using the Infrared Surface Brightness (IRSB) technique have uncertainty due to mass loss and infrared excess.  The infrared surface brightness technique uses an empirical fit of the dependence of surface brightness on $(V-K)_0$ to determine the mean value and amplitude of a Cepheid's angular diameter  \\citep{Gieren1999}. This angular diameter information is then converted to a radius and a distance by using radial velocity observations to compute the amplitude of the radius variation.  This technique has been shown to be robust, but infrared excess from mass loss suggests that the color $(V-K)_0$ may be overestimated.  The objective of this paper is to analyze fundamental-mode Cepheids in the LMC that have stellar fluxes from the OGLE-III survey and infrared data from the 2MASS and from two epochs of the SAGE surveys, and test if the Cepheids have infrared excess.  We determine mass-loss rates from infrared excesses determined from a correlation of these data sets by assuming a mass-loss model.  In the next section we describe the method for testing the model.  In section 3 we fit the model assuming zero mass loss, and the results of fitting the mass-loss model are given in section 4.  In section 5, we test how the removal of the $5.8$ and $8.0$ $\\mu m$ data affect the mass-loss model as well as how the unknown pulsation phase of the IR fluxes affects the predicted mass-loss rates.  The statistical comparison of the two models and the uncertainty of the fits due to the unknown pulsation amplitudes of the IR fluxes is given in section 5.  Using the results of the models, we calculate the Period-Radius relation (section 6), and test the uncertainty of the IRSB technique due to mass loss (section 7). In Section 8 we compute linear IR P-L relations from the data and compare the predictions with the observed results of \\cite{Ngeow2008b}, \\cite{Freedman2008}, and \\cite{Madore2008} in Section 9.  In Section 10, we test the P-L relations for non-linearity. ",
        "conclusions": ""
    },
    "1005/1005.0552_arXiv.txt": {
        "abstract": "The Cherenkov radio pulse emitted by hadronic showers of energies in the EeV range in ice is calculated for the first time using  full three dimensional simulations of both shower development and the coherent radio pulse emitted as the excess charge develops in the shower. A Monte Carlo, \\zhaires, has been developed for this purpose combining the high energy hadronic interaction capabilities of AIRES, and the dense media propagation capabilities of TIERRAS, with the precise low energy tracking and specific algorithms developed to calculate the radio emission in ZHS. A thinning technique is implemented to allow the simulation of radio pulses induced by showers up to 10 EeV in ice. The code is validated comparing the results for electromagnetic and hadronic showers to those obtained with GEANT4 and ZHS codes. The contribution to the pulse of other shower particles in addition to electrons and positrons, mainly protons, pions and muons, is found to be below 3 $\\%$ for 10 PeV and above proton induced showers. The characteristics of hadronic showers and the corresponding Cherenkov frequency spectra are compared with those from purely electromagnetic showers. The dependence of the spectra on shower energy and high-energy hadronic model is addressed and parameterizations for the radio emission in hadronic showers in ice are given for practical applications. ",
        "introduction": "The observation of neutrinos of EeV-scale energies is one of the main priorities in Astroparticle Physics. The detection of neutrinos will open a new window to observe parts of the Universe shielded by large depths of matter, unaccessible to conventional astronomy. The measurement of diffuse neutrino fluxes will provide further clues to the identification of the sources of ultra-high energy cosmic ray production, their composition and their production mechanisms~\\cite{becker08}. In addition such detections could have important implications for fundamental particle physics.  A very promising and cost-effective method to detect high-energy neutrino interactions was first proposed by G.~A.~Askar'yan in the 1960's \\cite{Askaryan62}. The idea is to detect the Cherenkov radiation at radio wavelengths generated by the excess number of electrons in the cascade of particles resulting from a high-energy particle interaction in a dense medium transparent to radio waves. The development of this excess charge, due to the interactions with matter electrons, is often referred to as the Askar'yan effect. The effect has been experimentally confirmed in accelerator experiments at SLAC in media such as sand \\cite{Saltzberg_SLAC_sand,Miocinovic_SLAC_sand}, rock salt \\cite{Gorham_SLAC_salt} and ice \\cite{Gorham_SLAC_ice}, with results in good agreement with theoretical calculations \\cite{alvz06}. At radio frequencies (MHz-GHz), the emission is coherent and the radiated power scales with the square of the primary particle energy. This makes this method very promising for the detection of neutrinos and cosmic rays of the highest energies (EeV) \\cite{zheleznykh,ZHS91,frichter96}.  Several experiments have already exploited this technique searching for ultra-high energy neutrinos, but no positive detection has been reported so far. These include experiments using the ice cap at the South Pole, such as the ANITA balloon experiment \\cite{ANITAlite,ANITAlong,ANITA_2009_limits} and the RICE array of antennas buried under the Antarctic ice \\cite{RICE03,RICElimits}. Other arrays are beginning to be constructed, such as the Askar'yan Radio Array ARA \\cite{ARA} and ARIANNA \\cite{ARIANNA}. There are also a number of projects using the Moon as target for neutrino interactions, along with radio telescopes on Earth as radiation detectors, such as the pioneering Parkes \\cite{Parkes96,Parkes07}, GLUE \\cite{GLUElimits}, Kalyazin \\cite{Kalyazin}, NuMoon \\cite{NuMoon}, LUNASKA \\cite{LUNASKA} and RESUN \\cite{RESUN}.  The interpretation of data from these experiments requires a detailed knowledge of the magnitude, angular distribution and frequency-dependence of the emitted Cherenkov radiation, which can then be related back to the properties of the induced cascade. This calls for accurate simulations of the Fourier-spectrum of coherent Cherenkov radiation from EeV showers in different dense media. In the past, full simulations of electromagnetic (EM) showers in ice were done up to PeV energies, using the well-known and well-tested ZHS code \\cite{ZHS91,ZHS92} and different versions of the GEANT code \\cite{razzaque04,almvz03}. Also, full simulations of hadronic and neutrino-induced showers in ice were carried out up to 100 TeV with GEANT \\cite{McKay_radio}. Hybrid Monte Carlo \\cite{alvarez_hybrid} and thinning techniques \\cite{Hillas} were also developed to simulate electromagnetic \\cite{alz97,aljpz09,acorne07}, hadronic \\cite{alz98,acorne07} and neutrino-induced showers \\cite{acorne07,alvz99} above PeV energies (mainly in ice and water). Above these energies the LPM effect \\cite{LPM,Stanev_LPM} starts to be important in the media under experimental consideration \\cite{Stanev_LPM,RalstonLPM,alz97}. Semi-analytical calculations of the radio-emission have also been performed \\cite{alvz00,buniy02}. Very recently, simulations of not only the Fourier-components of the spectrum but also of the radio-pulse in the time-domain emitted in electromagnetic showers, have been performed with the ZHS code\\cite{alrwz10}.  Modeling hadronic showers up to EeV energies is of utmost importance for neutrino detection, since they are induced by all neutrino flavors in neutral current (NC) interactions, as well as in the hadronic vertex of charged current (CC) interactions of muon and tau neutrinos. Although at EeV energies only $\\sim 20\\%$ of the neutrino energy is on average carried by the hadronic shower, these showers are known to be less affected by the LPM effect \\cite{alz98}, and their Cherenkov emission does not suffer from the shrinking of the Cherenkov cone, which would otherwise reduce the solid angle for observation. On the other hand, mixed showers induced in CC electron neutrino interactions are composed of a purely electromagnetic shower, produced by an energetic electron carrying on average $\\sim 80\\%$ of the energy of the neutrino, and a hadronic shower, initiated by the debris of the interacting nucleon. The advantage for neutrino detection of these types of interaction is that all the neutrino energy is channeled into the resulting shower. However, the observation of the electromagnetic shower is expected to be very difficult, since the LPM effect shrinks the angular distribution of the electric field, reducing dramatically the available solid angle for detection. For these reasons, experiments aiming at detecting neutrinos using the radio technique gain most of their acceptance from neutrino-induced hadronic showers \\cite{RICElimits,ANITAlong,alvarez_icrc09}.  Clearly, accurate simulations of hadronic showers with energies above which the LPM effect becomes effective in different dense media are needed. In this paper we present the first steps in that direction. We have used the well-known AIRES code \\cite{aires} in combination with the TIERRAS package \\cite{TIERRAS} to simulate proton-induced showers in dense media, such as ice. We have implemented the algorithms to calculate the Fourier components of the electric field produced by charged particle tracks in the shower, developed by Zas-Halzen-Stanev in the well-known and well-tested ZHS code \\cite{ZHS91,ZHS92}. The result is a flexible and powerful code named \\zhaires, which allows the simulation of electromagnetic, hadronic and neutrino-induced showers in a variety of media, along with their associated coherent radio emission due to the Askar'yan effect.   It is important to stress that in \\cite{alz97,alz98} the calculations of the emission properties of showers in ice up to EeV energies were only done in an approximate way. They were based on the Fourier-transform of simulated longitudinal and lateral shower profiles. With \\ZHAireS we can obtain the radio-pulse features in a consistent manner within a well-tested simulation, computing the electric field emitted by charged particles in the shower. In Ref.\\cite{McKay_radio}, GEANT simulations were used to calculate the radio-pulse properties in the same way as the \\ZHAireS code. However these calculations were limited to energies below 100 TeV due to GEANT limitations, while with \\ZHAireS we can simulate showers at EeV energies and above. Also, and for the first time, the contribution to the radio pulse of charged pions, muons and protons is accounted for and quantified, and parameterizations of the frequency spectrum of the radio pulse due to Askar'yan effect in hadronic showers are presented. Previous parameterizations of hadronic showers exist in the literature~\\cite{alz98}, but these are based on less accurate simulations than those presented in this work.   This paper is structured as follows. In Section \\ref{newcode} we describe the new \\ZHAireS code and compare its performance and results for electromagnetic showers with those of the well-tested codes GEANT 4 \\cite{GEANT4} and ZHS \\cite{ZHS92,aljpz09}. We also explore thinning techniques in the \\ZHAireS code, essential for the simulation of EeV showers. Section \\ref{hadronic} is devoted to the simulation of hadronic showers up to EeV energies, emphasizing their differences with respect to purely electromagnetic showers. Section \\ref{conclusions} summarizes and concludes the paper. In Appendix A, we also give parameterizations of the intensity of the Fourier-spectrum of the Cherenkov electric field emitted in hadronic showers, as a function of shower energy and observation angle.  In this paper we concentrate on hadronic showers in ice. We defer to future papers the study of coherent radio emission from EM and neutrino-induced showers. However we show that our results for proton showers can be applied to approximately model the Cherenkov emission from hadronic showers induced in high-energy neutrino interactions. The code can also be applied to the simulation of radio emission in extensive air showers~\\cite{zhairesarena2010,zhairesair2011}, despite the different emission mechanism \\cite{allan71,ARENA08}. ",
        "conclusions": "\\label{conclusions}  We have presented \\zhaires, a Monte Carlo code that merges the high energy hadronic interaction and tracking capabilities of AIRES\\cite{aires}, and the dense media propagation capabilities of TIERRAS \\cite{TIERRAS} with the precise low energy $e^\\pm$ tracking and radiation calculation capabilities of ZHS\\cite{ZHS91,ZHS92}. This combination allows a precise full simulation of radio emission of electromagnetic, and for the first time also hadronic showers up to EeV energies. \\ZHAireS has been used in conjuction with the TIERRAS package \\cite{TIERRAS} to simulate showers in dense media, obtaining the radiation emitted due to the Askar'yan effect.  We have compared \\ZHAireS results for electromagnetic showers in ice against ZHS up to 10 EeV, obtaining a good agreement between them. In the case of hadronic showers, we have compared the results from \\ZHAireS and GEANT4 up to 90 TeV \\cite{McKay_radio}, the highest energy for full simulations of hadronic showers with GEANT in the literature, also with very good quantitative agreement. Unlike GEANT4, \\ZHAireS is capable of simulating the emission from hadronic showers up to EeV energies.  By comparing the results of thinned \\ZHAireS simulations with un-thinned ones up to 1 PeV, we obtained thinning parameters which give a good compromise between accuracy and CPU time. We then tested the applicability of these parameters at higher energies, by comparing our results against ZHS, which has its own, well established \\cite{aljpz09}, thinning algorithms and parameters.  We confirmed the expectation that the EM energy ratio of hadronic showers keeps approaching that of EM showers at the highest energies, increasing the magnitude of the electric field (normalized by $E_0$) as energy grows. We also showed that, in the case of UHE purely EM showers, the EM energy ratio deviates from a linear dependence on energy, due to photohadronic interactions.  By comparing EM and hadronic showers up to 10 EeV we have found that, at the Cherenkov angle, the cut-off frequency tends to be smaller in hadronic showers and increases slowly as the energy rises because the lateral distribution becomes narrower. For angles away from the Cherenkov angle, as expected, the cut-off frequency decreases with energy, due to the logarithmic growth of the longitudinal profile. In contrast, away from the Cherenkov angle, the cut-off frequency for EM showers decreases rapidly above PeV energies due to the LPM effect, which is much more pronounced in EM showers than in hadronic ones. We also found that the fluctuations in hadronic showers are almost independent of energy, while in EM showers they grow rapidly, as the LPM effect becomes important.  We have also analyzed the influence of charged pions, muons and protons on the radio emission of showers. We found that most of their contribution is due to the excess of protons, with pions being the next contribution in importance, that induce a decrease of the (total) negative excess track length of the shower of $\\sim1-2\\%$ above 1 PeV. Although the actual field emitted by the hadrons in the shower is very small, the hadronic component carries a significant fraction of the shower energy. The importance of this energy balance is illustrated by the effects of the EM energy fraction on the radio emission. Given our results, following the actual tracks of hadrons in the shower turns out not to be essential for the field calculation, but accounting for the hadronic energy of the shower is.   We have compared the results obtained using SIBYLL 2.1 \\cite{SIBYLL21} and QGSJET01 \\cite{QGSJET01}, and found that there is a small dependence of the normalization of the emitted electric field on the hadronic model, up to $\\sim 4\\%$ lower in QGSJET01 at the Cherenkov angle. This is due to the higher number of protons, pions and muons when QGSJET01 is used, which increases the positive projected excess track-length of protons, pions and muons, with protons dominating the difference. Also more protons, pions and muons means means less electrons and positrons in the shower, decreasing the negative excess track. The net effect is a lower negative projected excess track as a whole, diminishing the field. At angles away from the Cherenkov angle, the longitudinal excess profile becomes important, and the shorter showers obtained using QGSJET lead to the maximum of the emission at higher frequencies, when compared to SIBYLL showers.  Finally, we have used our simulations of the emission due to the Askar'yan effect in hadronic showers to obtain new parameterizations of the radio pulse frequency spectrum which are described in Appendix A. We note that our results for proton showers can also be applied to approximately model the Cherenkov emission from hadronic showers induced in high-energy neutrino interactions. This is shown in the bottom panel of Fig.~\\ref{fig:MC_vs_param} where we have plotted the frequency spectrum obtained in \\ZHAireS simulations of a neutrino-induced shower in a neutral current interaction (or a charged current interaction of a muon neutrino) in which the secondaries from the fragmentation of the nucleon carry 2 EeV of energy in total. The products of the neutrino-nucleon interaction, mainly pions, were obtained with the HERWIG interaction Monte Carlo code \\cite{HERWIG}. The frequency spectrum obtained in \\ZHAireS simulations is compared to the parameterization of a proton shower with energy $E_0=2$ EeV as given in Appendix A."
    },
    "1005/1005.4767_arXiv.txt": {
        "abstract": "Type I X-ray bursts are thermonuclear stellar explosions driven by charged-particle reactions. In the regime for combined H/He-ignition, the main nuclear flow is dominated by the {\\it rp-process} (rapid proton-captures and $\\beta^+$-decays), the $3\\alpha$-reaction, and the {\\it $\\alpha$p-process} (a suite of ($\\alpha$,p) and (p,$\\gamma$) reactions). The main flow is expected to proceed away from the valley of stability, eventually reaching the proton drip-line beyond A = 38. Detailed analysis of the relevant reactions along the main path has only been scarcely addressed, mainly in the context of parameterized one-zone models. In this paper, we present a detailed study of the nucleosynthesis and nuclear processes powering type I X-ray bursts. The reported 11 bursts have been computed by means of a spherically symmetric (1D), Lagrangian, hydrodynamic code, linked to a nuclear reaction network that contains 325 isotopes (from $^1$H to $^{107}$Te), and 1392 nuclear processes. These evolutionary sequences, followed from the onset of accretion up to the explosion and expansion stages, have been performed for 2 different metallicities to explore the dependence between the extension of the main nuclear flow and the initial metal content. We carefully analyze the dominant reactions and the products of nucleosynthesis, together with the the physical parameters that determine the light curve (including recurrence times, ratios between persistent and burst luminosities, or the extent of the envelope expansion). Results are in qualitative agreement with the observed properties of some well-studied bursting sources. Leakage from the predicted SbSnTe-cycle cannot be discarded in some of our models. Production of $^{12}$C (and implications for the mechanism that powers superbursts), light p-nuclei, and the amount of H left over after the bursting episodes will also be discussed. ",
        "introduction": "Type I X-ray bursts (hereafter, XRBs) are cataclysmic stellar events. They are powered by thermonuclear runaways (TNR) in the H/He-rich envelopes accreted onto neutron stars in close binary systems (see reviews by Bildsten 1998; Lewin et al.  1993, 1995; Psaltis 2006; Schatz \\& Rehm 2006; Strohmayer \\& Bildsten 2006).  These events constitute the most frequent type of thermonuclear stellar explosion in the Galaxy (the third, in terms of total energy output after supernovae and classical novae), in part because of their short recurrence period (hours to days). About $90$ Galactic low-mass X-ray binaries exhibiting such bursting behavior (with burst durations of $\\tau_{burst} \\sim$ 10 - 100 s) have been found since the discovery of XRBs by Grindlay et al. (1976), and independently, by Belian et al. (1976). Type I X-ray bursts and their associated nucleosynthesis have been extensively modeled by different groups (see pioneering work by Woosley \\& Taam 1976, Maraschi \\& Cavaliere 1977, and Joss 1977), reflecting the astrophysical interest in determining the nuclear processes that power the explosion, the light curve, as well as in providing reliable estimates for the chemical composition of the neutron star surface (see Schatz et al. 1999, Parikh et al. 2008, and references therein).  With a neutron star hosting the explosion, temperatures and densities in the accreted envelope reach high values: T$_{peak} > 10^9$ K, and $\\rho \\sim 10^6$ g cm$^{-3}$. As a result, detailed nucleosynthesis studies require the use of hundreds of isotopes, linked by thousands of nuclear interactions, extending all the way up to the SnSbTe-mass region (Schatz et al. 2001) or beyond (the extent of the nuclear activity\\footnote{The nuclear activity reflects the changes in composition driven by different nuclear processes (i.e., p- and $\\alpha$-capture reactions, $\\beta$-decays, ...) that take place in the envelope at different stages of the burst. In this work, the extent of the nuclear activity is arbitrarily defined by the heaviest nucleus that achieves a mass fraction $>10^{-9}$.} in the XRB nucleosynthesis study of Koike et al. 2004 reaches $^{126}$Xe). Indeed, the extent of the rp-process in XRBs is still not clear: recent experimental work now shows that it will be more difficult to reach the SnSbTe-mass region (Elomaa et al. 2009). Because of computational constraints, XRB nucleosynthesis studies have been traditionally performed using limited nuclear reaction networks, truncated near Ni (Woosley \\& Weaver 1984; Taam et al. 1993; Taam, Woosley, \\& Lamb 1996 --all using a 19-isotope network), Kr (Hanawa, Sugimoto, \\& Hashimoto 1983 --274-isotope network; Koike et al. 1999 --463 nuclides), Cd (Wallace \\& Woosley 1984 --16-isotope network), or Y (Wallace \\& Woosley 1981 --250-isotope network). On the other hand, Schatz et al. (1999, 2001) have carried out very detailed nucleosynthesis calculations with a network containing more than 600 isotopes (up to Xe, in Schatz et al. 2001), but using a one-zone approach. Koike et al. (2004) have also performed detailed one-zone nucleosynthesis calculations, with temperature and density profiles obtained from a spherically symmetric evolutionary code, linked to a 1270-isotope network extending up to $^{198}$Bi.  Until recently, it has not been possible to couple hydrodynamic stellar calculations (in 1-D) and detailed networks. Recent efforts include Fisker et al. (2004, 2006, 2007, 2008), and Tan et al. (2007) ($\\sim$ 300 isotopes, up to $^{107}$Te), Jos\\'e \\& Moreno (2006) (2640 reactions and 478 isotopes, up to Te), or Woosley et al. (2004) and Heger et al. (2007) (up to 1300 isotopes with an adaptive network). This has prompted a detailed analysis of the nuclear activity powering the bursts. The most detailed work to date is that of Fisker et al. (2008), in the context of the 1-D general relativistic hydrodynamic code AGILE (Liebend\\\"orfer et al. 2002), linked to a nuclear reaction network containing 304 isotopes: a thorough analysis of the main nuclear activity in one characteristic burst is reported (although details for a sequence of 5 consecutive, 'representative' bursts are also outlined). However, because of the specific choice of metallicity (Z = $10^{-3}$, for the accreted matter) and mass-accretion rate ($\\dot M \\sim 10^{17}$ g s$^{-1}$) adopted, the nuclear activity does not extend much beyond mass $A \\sim 65$, as a result of compositional inertia effects, that quench further extension of the nuclear path. Hence, the flow does not reach the {\\it SnSbTe}-mass region, which was suggested as a natural endpoint in XRB nucleosynthesis studies (see Schatz et al. 1999,2001).  Clearly, the identification of the most relevant reactions in the $A \\sim 65 - 100$ mass region remains to be addressed in detail in the framework of hydrodynamic simulations. This is particularly relevant since, as first pointed out by Hanawa et al. (1983), proton captures on heavy nuclei (i.e., the {\\it rp}-process) have a dramatic effect on the shape of XRB light curves. To this end, a new set of type I X-ray bursts have been computed with SHIVA, a 1-D, spherically symmetric, hydrodynamic, implicit, Lagrangian code, used extensively in the modeling of classical nova outbursts (see Jos\\'e \\& Hernanz 1998). The code has been linked to a fully updated nuclear reaction network containing 324 nuclides and 1392 nuclear processes, a subset of that used in Parikh et al. (2008), and includes the most relevant charged-particle induced reactions occurring between $^1$H and $^{107}$Te, as well as their corresponding reverse processes. It is worth noting that the size of this network is similar (though slightly larger) to that adopted by Fisker et al. (2008). In order to set up the reaction rate library for our study, we started by adopting the proton drip line predicted by Audi et al. (2003a, 2003b). Experimental rates are available for a small subset of reactions (adopted from Angulo et al. 1999, Iliadis et al. 2001, and some recent updates for selected reactions). For all other reactions for which experimental rates are not available, we used the rates from the Hauser-Feshbach codes MOST (Goriely 1998; Arnould \\& Goriely 2006) and NON-SMOKER (Rauscher \\& Thielemann 2000; for details see Parikh et al. 2008). Neutron captures are disregarded since our early test calculations revealed that they play a minor role in XRB nucleosynthesis. All reaction rates incorporate the effects of thermal excitations in the target nuclei (Rauscher \\& Thielemann 2000). Screening factors are taken from Graboske et al. (1973) and DeWitt, Graboske, \\& Cooper (1973). For the weak interactions, $\\beta$-delayed nucleon emission and laboratory decay rates (Audi et al. 2003a) have been adopted. For a discussion of employing stellar versus laboratory decay rates, see Woosley et al. (2004). It is worth noting, however, that many computed stellar decay rates ( Fuller et al. 1982a, 1982b; Langanke \\& Martinez-Pinedo 2000) do not converge to their laboratory values at lower temperatures and densities, calling into question the model used for these calculations. Studies employing properly converging stellar decay rates for all isotopes relevant to XRB nucleosynthesis have not been performed by any group yet, and would certainly be interesting, although the results presented in this work would not be dramatically affected by their inclusion.  SHIVA uses a time-dependent formalism for convective transport whenever the characteristic convective timescale becomes larger that the integration time step. Partial mixing between adjacent convective shells is treated by means of a diffusion equation (Prialnik, Shara, \\& Shaviv 1979). No additional semiconvection or thermohaline mixing is considered. Models make use of Iben's (1975) opacity fits, better suited than the OPAL opacities for astrophysical environments that exhibit strong variations in metallicity, as in XRB nucleosynthesis. However, plans to incorporate these more realistic opacities are currently underway. The adopted equation of state includes contributions from the electron gas (with different degrees of degeneracy; Blinnikov et al. 1996), a multicomponent ion plasma, and radiation; Coulomb corrections to the electronic pressure are also taken into account.  Accretion is computed by redistributing material through a constant number of envelope shells (see Kutter \\& Sparks 1980, for details). To handle this, a tiny envelope, containing $1.1 \\times 10^{18}$ g of material (less than 1 permil of the total envelope mass accreted during the first bursting episode), distributed through all the envelope shells, is put initially in place(the influence of the number of envelope shells on burst properties will be discussed in Section 3). The model is then relaxed using a few, very large timesteps, to guarantee hydrostatic equilibrium. The temperature at the bottom of the envelope barely reaches $2.7 \\times 10^7$ K, whereas the density is just $1.4 \\times 10^3$ g cm$^{-3}$ (corresponding to a pressure of $5.7 \\times 10^{18}$ dyn cm$^{-2}$).  Mass accretion and nuclear reactions are then initiated.  Special emphasis is placed on the effect of the initial metallicity of the accreted matter on the main nuclear path, which in turn, will affect the final post-burst envelope composition and the shape of the light curves.  The structure of the manuscript is as follows: in Section 2, we analyze the main features (nuclear path, nucleosynthesis, light curves, etc) of a series of 4 bursts computed in a model with solar-like accreted material. The effect of the resolution adopted in this model is discussed in Section 3. A detailed analysis of the impact of the metallicity of the accreted material on burst properties is given in Section 4. Finally, a comparison with previous work, together with a thorough analysis of the corrections posed by general relativity, are discussed in Section 5. ",
        "conclusions": "\\subsection{General relativity corrections} The calculations reported here have been performed assuming Newtonian gravity. Since the envelope layers are very thin, it is easy to introduce general relativity corrections to this Newtonian framework (see Ayasli \\& Joss 1982, Lewin et al. 1993, Taam et al. 1993, Cumming et al. 2002, and Woosley et al. 2004). To this end, the surface gravity is rewritten as $g = G M_*/R_*^2 (1+z)$, where M$_*$ is the mass, R$_*$ is the stellar radius (defined in such a way that the surface area is 4$\\pi R_*^2$), and $z$ is the gravitational redshift given by $1 + z = (1 - 2G M_*/R_*c^2)^{-1/2}$. Our models of M$_* = 1.4$ M$_\\odot$ require R$_* = 14.3$ km, and a gravitational redshift of $z = 0.19$.  Following Woosley et al. (2004), once the redshift and radius are determined, it is straightforward to derive the set of correcting factors to the physical magnitudes described above for a suitable observer at infinity. Hence, recurrence times and burst durations should be increased by a factor $1 + z$. The mass-accretion rate as well as the burst luminosity have to take into account both the difference in surface area (compared to the Newtonian framework) and the gravitational redshift term. The energy and rest mass-accretion rate scale as R$_*^2$/$(1+z)$, while the luminosity $\\propto$ R$_*^2$/$(1+z)^2$. However, when M$_*$ is taken exactly as $M_{NS}$ (Newtonian framework), the surface area and redshift corrections for energy and mass accretion rate cancel out, since g $\\propto$ $(1+z)/$R$_*^2 = const$, and hence, no correction to the observed burst energy or mass-accretion rate is necessary, while the luminosity correction is simply given by $1/(1+z) = 0.84$. In addition, the accretion luminosity for an observer at infinity changes only by a factor 1.012, that is, the ratio between gravitational energy released per unit mass in general relativity, $c^2.z/(1+z)$, and the Newtonian value, $GM_{NS}/R_{NS}$. Finally, the luminosity measured at infinity will be smaller by a factor of $(1+z)=1.19$.  \\subsection{Comparison with previous work}  For consistency, the results discussed in this paper have been compared with those reported in previous work (obtained with similar hydrodynamic codes or in the framework of one-zone models).  As emphasized in Section 2, model 1 is qualitatively similar to model ZM of Woosley et al. (2004).  The twelve bursts computed by Woosley et al. (2004) in a Newtonian frame  were characterized by recurrence times of about $\\sim 2.7$ hr, peak luminosities of L$_{peak} \\sim (1.5 - 2) \\times 10^{38}$ erg s$^{-1}$, and ratios between persistent and burst luminosities of $\\alpha \\sim 60-65$. Our calculations (model 1, Newtonian frame) yield $\\tau_{rec} \\sim 5 - 6.5$ hr, L$_{peak} \\sim (3 - 7) \\times 10^{38}$ erg s$^{-1}$, and $\\alpha \\sim 35-40$.  The role played by the metallicity of the accreted material (model 3, with Z = Z$_\\odot$/20 = 0.001) qualitatively agrees with the pattern reported by Woosley et al. (2004) (see also, Heger et al. 2007). Longer recurrence times of $\\sim 9$ hours, peak temperatures of about $(1.3-1.4) \\times 10^{9}$ K, and ratios between persistent and burst luminosities of $\\alpha \\sim 20-30$ (with L$_{peak}$ $\\sim 10^{38}$ erg s$^{-1}$) have been obtained in the 5 bursts computed in model 3. In turn, the fifteen bursts computed by Woosley et al. (2004) for model zM are characterized by recurrence times of about $3 - 3.5$ hr, peak luminosities of L$_{peak} \\sim 10^{38}$ erg s$^{-1}$, and ratios between persistent and burst luminosities of  $\\alpha \\sim 50-60$. Results reveal a dependence of burst properties on the metallicity of the accreted material: the smaller the metal content, the larger the recurrence time (and the smaller the $\\alpha$). In turn, explosions in metal-deficient envelopes (i.e., model 3) are characterized by lower peak luminosities and longer decline times, in agreement with the pattern described in Woosley et al. (2004) and Heger et al. (2007). Model 3 bears as well a clear resemblance with the model computed by Fisker et al. (2008). In that work, five representative bursting sequences were analyzed, with $\\tau_{rec} \\sim 3.5 - 4$ hr, L$_{peak} \\sim (7 - 8) \\times 10^{37}$ erg s$^{-1}$, and $\\alpha \\sim 65-70$, as measured at infinity.  Despite the qualitative similaries in the gross properties of the bursts presented in this paper (as well as in the role played by the metallicity of the accreted material) and  those reported in previous work, a quantitative comparison reveals some discrepancies that are worth analyzing. In model 1 (with Z = Z$_\\odot$), our computations yield systematically larger (by a factor of $\\sim$ 2) recurrence times and peak luminosities (and hence, lower $\\alpha$) than model ZM of Woosley et al. (2004). Similar results are found in the low-metallicity case (model 3, with Z = Z$_\\odot$/20) when compared with model zM of Woosley et al. (2004), except for the peak luminosities that turn out to be very similar. It is also worth noting that the values reported by Fisker et al. (2008) show discrepancies with respect to Woosley et al. (2004), in particular, lower peak luminosities (and larger $\\alpha$). A major difference concerns the much larger effect played by the metallicity of the accreted material in this work as compared with Woosley et al. (2004), who explained the moderate effect found as due to compositional inertia washing out the influence of the initial metallicity. Another striking issue concerns the extremely large differences in the gross physical characteristics -nucleosynthesis, energies or recurrence times- between the first and subsequent bursts, as reported by Woosley et al. (2004). In terms of nucleoynthesis or nuclear activity, Figs. \\ref{fig:abun_b3}, \\ref{fig:lum_m3_5b}, \\& \\ref{fig:LUMI_OVER}, reveal a similar behavior for the different bursts (although a somewhat lower production of intermediate-mass elements as well as of the heaviest elements is reported for the first burst computed in model 3).  Very limited information on the nucleosynthetic yields obtained in model ZM is given in Woosley et al. (2004). Thus, we will restrict the discussion on the extent of the nuclear activity and on the resulting chemical abundance pattern to model 3, through a brief comparison with the work reported by Schatz et al. (2001), Fisker et al. (2008) and Woosley et al. (2004) (for model zM). It is worth noting that both the nucleosynthetic end-point (located in the SnSbTe-mass region) and the main nuclear path in the A $\\sim 50 - 100$ mass region obtained in this work (passing through a suite of different nuclei such, as $^{55}$Co, $^{60}$Zn, $^{70}$Br, $^{75}$Rb, $^{85}$Mo, $^{90}$Rh, or $^{100,105}$Sn) are similar to those reported by Schatz et al. (2001) in the framework of one-zone calculations. Whereas the main nuclear path in the first burst of model zM (Woosley et al. 2004) is very similar to the one reported in this work, compositional inertia causes a more limited extension of the nuclear activity in the successive bursts of Woosley et al.: hence, while the three most abundant nuclei at the bottom of the envelope are $^{106}$Sn and $^{104,106}$In at the end of the first burst, this switches to $^{64}$Zn, $^{68}$Se, and $^{32}$S (a similar trend is also reported by Fisker et al. 2008). In this work, the mass-averaged composition at the end of the first burst computed for model 3 (see Table 8) is dominated (aside from some residual H and $^{4}$He) by the presence of $^{104}$Pd (0.05, by mass) and $^{105}$Ag (0.08), while X($^{64}$Zn) $\\sim$ 0.04. But the peak at the end of the abundance distribution (see Fig. \\ref{fig:abun_b3}) increases with subsequent bursts up to a plateau value, which indicates that these heavy nuclei are still produced in similar quantities. This is very different to the results reported by Woosley et al. (2004). Indeed, at the end of the fifth burst, the abundance pattern, shows still a significant presence of heavy species (i.e., X($^{105}$Ag) $\\sim$ 0.1, X($^{104}$Pd) $\\sim$ 0.08, and X($^{94}$Tc) $\\sim$ 0.05), together with a simultaneous increase in the abundances of intermediate-mass elements, such as $^{60}$Ni (0.06), $^{64}$Zn (0.09), $^{68}$Ge (0.07), or $^{72}$Se (0.04). It is finally worth noting that, in agreement with all previous hydrodynamic studies, both models 1 and 3 yield very small post-burst abundances of $^{12}$C, below the threshold amount required to power superbursts.  Even though only a few bursts have been computed for these models, they already show a trend on the amount of $^{12}$C that may be expected after many more bursts.  Finally, it is also worth mentioning that large differences exist between the hydrodynamic simulations reported here (see also Woosley et al. 2004, and Fisker et al. 2008) and those based on one-zone models (i.e., Schatz et al. 1999, 2001) as regards the shape of the light curve accompanying the bursting episodes (the primary difference being the presence of a long-lasting plateau in the latter).  The origin of the discrepancies reported is not totally clear and would require additional hydrodynamic studies. Notice, however, that the local surface gravity of our model is somewhat smaller than that adopted in the abovementioned works: whereas a 10 km radius is assigned to the 1.4 M$_\\odot$ neutron star in Woosley et al. (2004) ($g = 1.86 \\times 10^{14}$ cm s$^{-2}$), the integration of the neutron star structure from the core to its surface, in hydrostatic equilibrium, yielded 13.1 km (14.3 km, after general relativity corrections are introduced; see Subsection 5.1), for our 1.4 M$_\\odot$ neutron star (corresponding to a surface gravity of $g = 1.08 \\times 10^{14}$ cm s$^{-2}$); in turn, the calculations reported by Fisker et al. (2008), in a general relativity framework, relied on a 11 km (1.4 M$_\\odot$) neutron star, for which $g = 1.53 \\times 10^{14}$ cm s$^{-2}$. Although XRB properties depend weakly upon the neutron star mass (or surface gravity), part of the differences outlined between the three studies can be attributed to the combined effect of the adopted neutron star size (surface gravity) and to differences in the input physics (i.e., nuclear reaction network, opacities, treatment of convection). In particular, the use of Iben's opacities may have some effect on the peak luminosities achieved since the larger OPAL opacities will likely decrease the amount of energy radiated away from the star. Moreover, the inclusion of semiconvection and thermohaline mixing would have a minor effect in the properties of the explosions, likely affecting the appearance of marginal convective transport between bursts (see Woosley et al. 2004, Fisker et al. 2008). It is however worth noting that the convective pattern shown in Figs. \\ref{fig:conv_tem_b4} \\& \\ref{fig:LUMI_OVER} is similar to those reported in previous work: namely that convection sets in as soon as superadiabatic gradients are established in the envelope, following the early stages of the TNR and the corresponding rise in temperature; it reaches the surface and begins to recede before the observed burst properly commences, shutting off thereafter (Woosley et al. 2004, Fisker et al. 2008).  The potential impact of XRB nucleosynthesis on Galactic abundances is still a matter of debate. Matter accreted onto a neutron star of mass $M$ and radius $R$ releases $G M m_p/R \\sim 200$ MeV nucleon$^{-1}$, whereas only a few MeV nucleon$^{-1}$ are released from thermonuclear fusion. Thus ejection from a neutron star is unlikely. However, it has been suggested that radiation-driven winds during photospheric radius expansion may lead to ejection of a tiny fraction of the envelope (containing nuclear processed material; see Weinberg et al. 2006; MacAlpine et al. 2007). Indeed, XRBs have been proposed as a possible source of the light p-nuclei $^{92,94}$Mo and $^{96,98}$Ru (Schatz et al. 1998, 2001). No matter is ejected in any of the models reported in this work, a result fully independent of the adopted resolution and in agreement with all previous hydrodynamic simulations (Woosley et al. 2004, Fisker et al. 2008). Moreover, it is worth noting that, as shown in Figs. \\ref{fig:abun_b1} \\& \\ref{fig:abun_b3}, the abundances of many species synthesized during the bursts decrease remarkably towards the outer envelope layers, because of inefficient convective transport (see Figs. \\ref{fig:conv_tem_b4} \\& \\ref{fig:LUMI_OVER}). To assess the possible contribution to the Galactic abundances, one has to rely on the abundances of the outer envelope layers (the only ones that have a chance to be ejected by radiation-driven winds). This shows the limitations posed by one-zone nucleosynthesis calculations, in which the chemical species synthesized in the innermost layers are, by construction, assumed to represent the whole (chemically homogeneous) envelope. The mass fractions of these p-nuclei, obtained in model 3, drop by more than an order of magnitude in the outer envelope layers (as compared with the values achieved at the innermost envelope; see Fig. \\ref{fig:MORU}); the resulting overproduction factors, $f \\sim 10^6$, are several orders of magnitude smaller than those required to account for the origin of these problematic nuclei (see Weinberg et al. 2006, Bazin et al. 2008), in sharp contrast with the results obtained on the basis of one-zone calculations (Schatz et al. 1998, 2001)."
    },
    "1005/1005.3830_arXiv.txt": {
        "abstract": "{Planets are formed in disks around young stars.  With an age of $\\sim$~10\\,Myr, \\object{TW~Hya} is one of the nearest T~Tauri stars that is still surrounded by a relatively massive disk. In addition a large number of molecules has been found in the \\object{TW~Hya} disk, making \\object{TW~Hya} the perfect test case in a large survey of disks with {\\it Herschel}--{\\it PACS} to directly study their gaseous component. We aim to constrain the gas and dust mass of the circumstellar disk around \\object{TW~Hya}. We observed the fine-structure lines of [\\ion{O}{I}] and [\\ion{C}{II}] as part of the Open-time large program {\\it GASPS}. We complement this with continuum data and ground-based $^{12}$ CO 3--2 and $^{13}$CO 3--2 observations.  We simultaneously model the continuum and the line fluxes with the 3D Monte-Carlo code {\\it MCFOST} and the thermo-chemical code {\\it ProDiMo} to derive the gas and dust masses. We detect the [\\ion{O}{I}] line at 63 $\\mu$m.  The other lines that were observed, [\\ion{O}{I}] at 145~$\\mu$m and [\\ion{C}{II}] at 157~$\\mu$m, are not detected. No extended emission has been found. Preliminary modeling of the photometric and line data assuming [$^{12}$CO]/[$^{13}$CO]=69 suggests a dust mass for grains with radius $<$ 1~mm of $\\sim$ 1.9 $\\times$ 10$^{-4}$ M$_{\\odot}$ (total solid mass of 3 $\\times$ 10$^{-3}$ M$_{\\odot}$) and a gas mass of (0.5--5) $\\times$ 10$^{-3}$ M$_{\\odot}$. The gas-to-dust mass may be lower than the standard interstellar value of 100.} ",
        "introduction": "Planets are formed in the disks that surround a large fraction of T~Tauri stars. Knowledge of the gas mass available at different disk ages is essential to constrain giant planet formation models. Most studies estimate the dust mass from millimeter continuum emission and assume the gas mass is a factor of 100 times larger.  This conversion factor has been calibrated for the interstellar medium but is likely not valid for disks, especially those that are evolving toward debris disks or where most of the gas has accreted onto the planetary atmosphere. Disk gas mass estimates derived from observations of $^{12}$CO and optically thinner $^{13}$CO emission are at least a factor of 10 lower than the mass derived from dust observations assuming the interstellar medium conversion factor.  The discrepancy has been ascribed to CO photodissociation at disk atmosphere and freeze-out onto cold dust grains in the disk midplane \\citep[e.g.,][]{Qi2004ApJ...616L..11Q,Thi2001ApJ...561.1074T}. An alternative explanation is that the CO abundance is not different and the gas in disks has been depleted.  The {\\it PACS} instrument \\citep{Poglistch2010} on-board the {\\it Herschel Space Telescope} \\citep{Pilbratt2010} makes it possible to observe lines from species that result from the photodissociation of CO (atomic oxygen and singly ionized carbon). With observations of all the major gas-phase carbon and oxygen-bearing species, we can more precisely constrain the disk gas mass.  At a distance of $\\sim$~56~pc \\citep{Wichmann1998MNRAS.301L..39W}, \\object{TW~Hya} is one of the nearest classical T~Tauri stars with an estimated age of 10~Myr \\citep{Barrado2006A&A...459..511B}.  Its proximity allows us to attain an order of magnitude higher mass sensitivity than objects in the Taurus molecular cloud.  Fits to the spectral energy distribution (SED) provide an estimate of the gas disk mass of 6 $\\times$ 10$^{-2}$ M$_{\\odot}$ after applying a conversion factor of $\\sim$~75 \\citep{Calvet2002ApJ...568.1008C}. This large disk mass at this advanced age is surprising as the median disk lifetime is only 2-3 Myr \\citep{Haisch2001ApJ...553L.153H}. \\object{TW~Hya} is considered a transition object with an optically thin inner cavity and an optically thick outer disk \\citep{Calvet2002ApJ...568.1008C,Ratzka2007A&A...471..173R}. The fit to the SED also suggests that grains have grown to at least $\\sim$~1~cm.  The star \\object{TW~Hya} was observed as a Science Demonstration Project object and is part of the {\\it Herschel-GASPS} program \\citep{Dent2010}. {\\it Herschel} observations of the disk around the Herbig~Ae star \\object{HD169142} are presented by \\citet{Meeus2010}. In this letter we use fine-structure lines in addition to continuum data and CO (sub)millimeter lines to directly constrain the gas mass and compare it to the dust mass derived from fits to the SED.   \\begin{figure}[!ht] \\centering {\\includegraphics[angle=90,width=8cm]{14578fig1.ps}} \\caption{{\\it Herschel-PACS} spectrum centred around the OI 63 $\\mu$m line on the upper-left panel.} \\label{fig_OI63_herschel} \\end{figure}   \\begin{table}% \\centering \\small \\caption{Lines observed by {\\it Herschel-PACS}. The errors and upper limits are 3~$\\sigma$. The calibration error adds an extra $\\sim$~40\\% uncertainty. The CO data also have uncertainties of 30\\%.} \\label{table_results} \\begin{tabular}{lllll} \\hline \\noalign{\\smallskip} \\multicolumn{1}{c}{Line} & \\multicolumn{1}{c}{Cont. flux} & \\multicolumn{1}{c}{Obs.}&\\multicolumn{1}{c}{GH08}&\\multicolumn{1}{c}{M08}\\\\ &  \\multicolumn{1}{c}{(Jy)}& \\multicolumn{3}{c}{(10$^{-18}$ W m$^{-2}$)}\\\\  \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} O{\\sc I}\\ $^3$P$_1 \\rightarrow ^3$P$_2$ & 2.99~$\\pm$~0.14 & \\phantom{$<$}36.5~$\\pm$~12.1 & 124-161 & 412\\\\ O{\\sc I}\\ $^3$P$_0 \\rightarrow ^3$P$_1$ & 7.00~$\\pm$~0.05 &       $<$\\  \\ 5.5 & 25-41 & 11\\\\ C{\\sc II}\\ $^2$P$_{3/2} \\rightarrow ^2$P$_{1/2}$ & 8.79~$\\pm$~0.08  &       $<$\\  \\ 6.6& 0.8-12 & 0.06\\\\ CO 3--2 & \\multicolumn{1}{c}{n.a.} & \\phantom{$<$}\\ \\ 0.43 & 0.3-0.6 & n.a.\\\\ $^{13}$CO 3--2 & \\multicolumn{1}{c}{n.a.} & \\phantom{$<$}\\ \\ 4.4 $\\times$ 10$^{-2}$& n.a.& n.a.\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table} \\normalsize ",
        "conclusions": "The {\\it Herschel-PACS} spectral observations were used to constrain the gas disk mass surrounding the 10 Myr T~Tauri star \\object{TW~Hya}. We estimate the gas mass to be (0.5--5) $\\times$ 10$^{-3}$ M$_\\odot$ compared to the dust mass ($a_{\\mathrm{max}}<$~1mm) of 1.9 $\\times$ 10$^{-4}$ M$_\\odot$. The gas-to-dust mass ratio is $\\sim$2.6--26, lower than the standard interstellar value of 100. The ratio gas-to-total-mass in solids is $\\sim$0.17--1.7.  Although the disk is still massive, a significant fraction of the primordial gas has already disappeared. A large fraction of the primordial gas may have been evaporated due to the strong X-ray flux from \\object{TW~Hya}. \\object{TW~Hya} is the first example where the disk gas mass around a transitional T~Tauri star can be determined accurately and directly from gas phase lines. However, more detailed modeling that includes X-ray physics and $^{13}$CO photochemistry is needed to confirm the low gas mass."
    },
    "1005/1005.2188_arXiv.txt": {
        "abstract": "We present a complete numerical study of cosmological models with a time dependent coupling between the dark energy component driving the present accelerated expansion of the Universe and the Cold Dark Matter (CDM) fluid. Depending on the functional form of the coupling strength, these models show a range of possible intermediate behaviors between the standard $\\Lambda $CDM background evolution and the widely studied case of interacting dark energy models with a constant coupling. These different background evolutions play a crucial role in the growth of cosmic structures, and determine strikingly different effects of the coupling on the internal dynamics of nonlinear objects. By means of a suitable modification of the cosmological N-body code {\\small GADGET-2} we have performed a series of high-resolution N-body simulations of structure formation in the context of interacting dark energy models with variable couplings. Depending on the type of background evolution, the halo density profiles are found to be either less or more concentrated with respect to $\\Lambda $CDM, contrarily to what happens for constant coupling models where concentrations can only decrease. However, for some specific choice of the interaction function the reduction of halo concentrations can be larger than in constant coupling scenarios. We also find that different types of coupling evolution determine specific features in the growth of large scale structures, like peculiar distortions of the matter power spectrum shape or different time evolutions of the halo mass function. Furthermore, also for time dependent couplings baryons and CDM develop a bias already on large scales, which is progressively enhanced for smaller and smaller scales, and the effect can be significantly larger compared to constant coupling scenarios. The same happens to the baryon fraction of halos, which can be more significantly reduced below its universal value in variable coupling models with respect to constant coupling cosmologies. In general, we find that time dependent interactions between dark energy and CDM can in some cases determine stronger effects on structure formation as compared to the constant coupling case, with a significantly weaker impact on the background evolution of the Universe, and might therefore provide a more viable possibility to alleviate the tensions between observations and the $\\Lambda $CDM model on small scales than the constant coupling scenario. ",
        "introduction": "\\label{i}  According to the present interpretation of the vast amount of cosmological data that have been collected over the last years, the Universe in which we live has a nearly flat spatial geometry, it expands with a rate of $\\sim 70 $ km s$^{-1}$ Mpc$^{-1}$, and the total amount of matter it contains accounts only for $\\sim 27\\%$ of the critical energy density that is required to justify its spatial flatness. Furthermore, the expansion seems to have entered an accelerated phase since about $6$ billion years, and the missing $\\sim 73\\%$ of the critical energy density is assumed to be in the form of some dark energy (DE) component able to drive such accelerated expansion. This detailed picture can be obtained by combining a wide variety of different and complementary cosmological datasets, ranging from Cosmic Microwave Backround (CMB) \\citep{wmap5,wmap7}, to Large Scale Structure surveys \\citep{Percival_etal_2001,Cole_etal_2005,Reid_etal_2010}, to observations of Type Ia Supernovae (SnIa) \\citep{Riess_etal_1998,Perlmutter_etal_1999,SNLS,Kowalski_etal_2008} and Baryon Acoustic Oscillations (BAO) \\citep[\\eg][]{Percival_etal_2009}. While the matter fraction of the Universe is known to be composed only for $\\sim 15\\%$ by baryonic particles, with the remaining mass in the form of some collisionless, weakly interacting Cold Dark Matter (CDM) component, the nature of the DE fraction is yet completely unknown, and its understanding constitutes one of the major challenges in modern cosmology.  The simplest possibility of a cosmological constant $\\Lambda $, whose energy density remains constant throughout the whole expansion history of the Universe,  is in good agreement with a very large amount of cosmological and astrophysical data, and this is the reason for the establishment of the $\\Lambda $CDM scenario as the present standard cosmological model. Nevertheless, the nature of the cosmological constant raises two fundamental questions concerning the very finely tuned value of its energy density (the ``fine tuning problem\") and the beginning of its domination over CDM only at relatively recent cosmological epochs (the ``coincidence problem\"). For this reason, alternative possibilities of dynamically evolving DE components have been proposed, in particular models where the DE is identified with a classical scalar field as for the case of quintessence \\citep{Wetterich_1988,Ratra_Peebles_1988} or k-essence \\citep{ArmendarizPicon_etal_2000,kessence}.  As a further extension of these dynamical DE scenarios, different possible forms of interaction between the DE component and the matter sector of the Universe have been suggested and investigated in the literature, as \\eg the generalized Chaplygin gas \\citep[see \\eg][]{Kamenshchik_etal_2001, Bilic_Tupper_Viollier_2002,Bento_etal_2002,Carturan_Finelli_2003,Amendola_etal_2003b}, unified dark matter models \\citep{Mainini_Bonometto_2004,Bertacca_etal_2007,Bertacca_Bartolo_Matarrese_2010}, extended quintessence models \\citep{Perrotta_etal_2000,Baccigalupi_Matarrese_2000,Pettorino_etal_2005}, and coupled DE \\citep{Wetterich_1995,Amendola_2000,Amendola_2004,Pettorino_Baccigalupi_2008}. Any form of interaction between the DE sector and other matter species, like CDM or massive neutrinos \\citep[as in \\eg ][]{Amendola_Baldi_Wetterich_2008} would leave distinctive features in the background expansion history of the Universe and in the growth of cosmic structures \\citep[see \\eg ][]{Brax_etal_2010} which could provide new ways to tackle the problem of the nature of DE. It is therefore essential to understand the impact that the interaction would have on observable quantities as \\eg the properties of CMB \\citep{Amendola_etal_2003,Amendola_Quercellini_2003,Bean_etal_2008,LaVacca_etal_2009}, of large scale structure formation \\citep{Koivisto_2005,Bean_etal_2008,Baldi_Pettorino_2010,Baldi_Viel_2010}, and of the nonlinear newtonian dynamics at small scales \\citep{Perrotta_etal_2003,Mainini:2006zj,Maccio_etal_2004,Saracco_etal_2010,Baldi_etal_2010}.  In the present work, in particular, we consider the case of a cosmological scenario where the DE scalar field interacts with the CDM fluid with a coupling strength that evolves in time, therefore generalizing the very widely studied case of constant couplings for the DE interaction. Other forms of effectively time dependent couplings, where the interaction depends on linear or nonlinear combinations of the energy densities of the interacting fluids and of their time derivatives have been studeid in \\eg \\citet{Barrow_Clifton_2006,CalderaCabral_2009,Chimento_2010}. Here we want to focus on general functional forms of the coupling strength in the context of coupled quintessence models where the interaction term is proportional to the energy density of the matter coupled fluid. One of the main motivations behind the introduction of a time dependence of the DE coupling to the matter sectors -- besides the fact that an evolving coupling is {\\em per se} a more general and natural assumption than a constant interaction strength -- lies in the recent discovery \\citep{Baldi_etal_2010} that the effects of the DE-CDM interaction on the formation and evolution of structures at small scales, in particular in the nonlinear regime, might help alleviating the tensions between the $\\Lambda $CDM model and a series of astrophysical observations. These range from  \\eg the abundance of satellites in CDM halos \\citep{Navarro_Frenk_White_1995}, to the observed low baryon fraction in large galaxy clusters \\citep{Ettori_2003,Allen_etal_2004,Vikhlinin_etal_2006,LaRoque_etal_2006,McCarthy_etal_2007}, to the so called ``cusp-core\" problem for the density profiles of the CDM halos of dwarf galaxies \\citep{Moore_1994,Flores_Primack_1994,Simon_etal_2003}, of spiral galaxies \\citep{Navarro_steinmetz_2000,Salucci_Burkert_2000,Salucci_2000,Binney_Evans_2001}, and of galaxy clusters \\citep{Sand_etal_2002,Sand_etal_2004,Newman_etal_2009}, or to the so called ``Dark Flow\" problem \\citep{Watkins_etal_2008}. Furthermore, some new potential challenges to the $\\Lambda $CDM scenario have been recently reported based on the detection of very massive high-redshift clusters \\citep[see \\eg][]{Jee_etal_2009,Rosati_etal_2009} or on the observed dynamical properties of CDM halo satellites \\citep{Lee_Komatsu_2010,Lee_2010}.  In their recent study \\citet{Baldi_etal_2010} showed by means of a series of high-resolution N-body simulations how the DE-CDM interaction -- for the case of constant coupling -- could alleviate some of these problems; in particular, it was shown how the interaction can reduce the ``cuspyness\" of massive CDM halos, thereby going in the right direction for a solution of the ``cusp-core\" problem. Nevertheless, the magnitude of this effect is strongly limited by the tight observational constraints on constant coupling models \\citep[as \\eg][]{Bean_etal_2008,LaVacca_etal_2009} that put a firm bound to the maximum allowed value of the coupling. It is therefore natural to speculate about the possibility that a time dependent coupling with a large value during the late stages of structure formation and a progressively smaller value at high redshift might increase significantly the impact of the new physics introduced by the interaction on \\eg the density profiles of CDM halos without perturbing the overall evolution of the Universe beyond the present observational limits.  The present work constitutes the natural extension of the analysis done by \\citet{Baldi_etal_2010} to the case of time dependent couplings. In this paper we perform a complete numerical study of some quite general classes of coupling functions, starting from their background evolution up to the nonlinear regime of structure formation, and we present the first high-resolution hydrodynamical N-body simulations of structure formation in the context of interacting DE models with a time dependent coupling to date.  The manuscript is organized as follows. In Section~\\ref{cde} we describe the main features of the coupled DE models under investigation with a particular stress on the differences with respect to the standard constant coupling models that have been widely studied in the literature. More specifically, in Sec.~\\ref{bkg} we discuss the background equations and in Sec.~\\ref{integration} we illustrate the numerical methods used to integrate such equations backwards in time. In Sec.~\\ref{obs} we discuss observational constraints on the coupled DE scenario and a possible way to use the bounds derived for constant coupling models to check the viability of the variable coupling cosmologies investigated in the present work. In Sec.~\\ref{prt} we study linear perturbations equations and we discuss the evolution of linear matter density fluctuations in variable coupling models.\\\\ In Sec.~\\ref{sim} we briefly summarize the numerical methods used in the N-body simulations, and in Sec.~\\ref{results} we present and discuss the results of our runs. Finally, in Sec.~\\ref{concl} we draw our conclusions. ",
        "conclusions": "\\label{concl}  In the context of interacting DE models we have studied a few general classes of time evolution of the interaction strength between DE and CDM, generalizing the widely studied case of constant couplings to the more natural scenario of a variable coupling.  Following the idea -- already discussed in previous works -- that a large value of the coupling could leave distinctive features in the properties of observable structures and even possibly alleviate the tensions between the $\\Lambda $CDM cosmological model and astrophysical observations at small scales, we have designed a few general forms of coupling functions $\\beta _{c}(\\phi )$ that grow in time, thereby having a significantly weaker impact on the overall background expansion of the Universe as compared to constant coupling models with the same interaction strength at $z=0$.  In particular, we have investigated three classes of time evolution of the coupling, where the interaction strength is proportional either to a power of the cosmological scale factor $a(t)$, or to the fractional DE density $\\Omega _{\\phi }$, or to an exponential function of the DE scalar field $\\phi $. The first two classes are purely phenomenological parametrizations of the time evolution of the coupling, while the latter one represents a more physical situation where the interaction depends on the dynamical evolution of the scalar field.\\\\  We have performed a complete numerical analysis of the background evolution for these three different types of coupling functions by solving the full system of coupled dynamic equations in the presence of a variable coupling, generalizing previous works. Even in the absence of analytic solutions for variable coupling models, our numerical integrations allow to identify the main background features of these cosmologies.  More specifically, we have shown that the first two phenomenological parametrizations of the coupling evolution mentioned above, due to the very fast decrease of the coupling with increasing redshift, do not present the so called ``$\\phi $MDE\" scaling solution typical of constant coupling models, and for what concerns their background evolution are practically indistinguishable from $\\Lambda $CDM, thereby suffering of the same level of fine tuning of the cosmological constant. On the contrary, the more physical form of a coupling that depends on the evolution of the scalar field shows a background evolution with an intermediate behavior between $\\Lambda $CDM and a standard ``$\\phi $MDE\" phase which can be still well reproduced by the usual analytic solution for the fractional DE density $\\Omega _{\\phi }$ during the ``$\\phi $MDE\" phase, once generalized to the case of growing couplings. We have therefore called this intermediate type of background evolution a ``Growing $\\phi $MDE\" phase.\\\\  We have then studied the evolution of linear perturbations within variable coupling models, pointing out the main differences arising in the perturbations equations due to the time dependence of the coupling. In particular, we have shown how in general a growing coupling function could induce instabilities in the growth of scalar perturbations at large scales, due to the presence of a negative effective mass term in the linear scalar field equation, and we have discussed under which conditions these instabilities can be avoided. We have then numerically computed the growth factor of matter density perturbations at subhorizon scales for a few selected models. Among these, the computed growth factors again clearly shows two distinct classes: on one side the phenomenological couplings present no significant differences in the growth of density perturbations with respect to $\\Lambda $CDM, except at very low redshifts, while on the other side the exponential coupling models have a faster growth of density fluctuations during most of the expansion history of the Universe.\\\\  The main focus of the present paper is on the effects of variable couplings on nonlinear structure formation. Exploiting the implementation of coupled DE cosmologies into the N-body code {\\small GADGET-2} developed for previous works, we have run high-resolution hydrodynamical N-body simulations for some selected cosmological models belonging to different classes of coupling evolution. For all these cosmologies initial conditions have been generated discarding possible early effects of the coupling based on the consideration that these effects were shown by previous studies to have a minor impact on the final properties of nonlinear structures at the present level of numerical resolution. We have also chosen to normalize all the cosmologies to the same amplitude of the large-scale power at the present time. Although this is a common choice, other conventions are equally valid and could lead to different predictions for the same cosmological models.\\\\  We have shown that the power spectrum of matter density fluctuations evolves in a strikingly different way in the two different types of models with respect to $\\Lambda $CDM. In particular, the phenomenological couplings have a very similar evolution to $\\Lambda $CDM until $z\\sim 0.5$, followed by a very fast growth of the power spectrum amplitude at scales below $k \\sim 1.0~h~\\text{Mpc}^{-1}$. On the contrary, the exponential models have a lower amplitude than $\\Lambda $CDM at all scales during most of the cosmic evolution, and catch up with the $\\Lambda $CDM power spectrum only at large scales at $z=0$, while at small scales they still show a significant lack of power also at the present time.  We have then confirmed that also for variable couplings the scalar fifth-force acting only between CDM particles induces a bias in the amplitude of density fluctuations in baryons and CDM at all scales, as it happens for constant coupling models. The bias shows a clear scale dependence that develops very quickly as the couplings approaches their large values at low redshift, and the enhancement of this effect when moving from the linear regime of very large scales to smaller and progressively more nonlinear scales is found to be stronger than for the constant coupling models studied in previous works.  This bias can be detected also in the very nonlinear regime characterizing the inner parts of collapsed objects, and has an impact on the total amount of baryons contained in massive halos that could therefore influence the determination of the baryon fraction from cluster measurements. We have therefore computed the evolution of the average baryon fraction within the virial radius $r_{200}$ for all the halos arising in the simulations of the different cosmological models, finding a generally stronger reduction of the baryon fraction as compared to constant couplings, with a decrease up to $\\sim 14 - 16 \\%$ with respect to $\\Lambda $CDM.  We have computed the mass functions at different redshifts for all of our selected models, showing how at $z=0$ all the cosmologies have similar shapes and amplitudes of the mass function, with relative differences of the order of $\\sim 10\\%$. However, at higher redshifts the mass functions of the exponential coupling models show a clear excess of small halos and a strong lack of large halos with respect to $\\Lambda $CDM, consistently with the later onset of structure formation in these models. We have also computed the multiplicity function for all the models and compared it with the theoretical predictions according to the \\citet{Sheth_Tormen_1999} and \\citet{Jenkins_etal_2000} fitting formulae, evaluated with the appropriate growth factor for all the models, and with the standard value of the extrapolated linear overdenisty at collapse. We found good agreement between the multiplicity functions in our simulated cosmologies and the theoretical fitting functions, with the only exception of one of the scale factor dependent models at $z=0$. This discrepancy might suggest the need to reconsider the spherical collapse formalism in the presence of strongly variable couplings, will be investigated in future works. Nevertheless, the usual mass function fitting formulae were found to be fairly accurate also for most of our variable coupling models, generalizing previous results.  Finally, we have investigated the effects of variable couplings on the halo density profiles. As for the case of constant coupling models, we find that they are still remarkably well fit in all the different cosmologies by the NFW formula. However -- in contrast with what happens within constant coupling models -- we also find that variable coupling cosmologies do not always show a decrease of the inner overdensity of halos with respect to the standard $\\Lambda $CDM case, but present opposite trends for the two different classes of coupling functions. The more realistic and physically motivated exponential coupling models show a significant decrease of the inner overdensity of halos with respect to $\\Lambda $CDM, while the phenomenological models show on the contrary a clear increase of the density in the central regions. This strikingly different behavior can be explained by considering which are the main physical mechanisms that can account for a modification of the equilibrium state of a collapsed halo in the context of our variable coupling models. As it was shown before for the case of constant couplings, the  mass decrease and the extra friction term in the equation of motion of CDM particles can only induce an increase of the total energy of a virialized system, which therefore restores its virial equilibrium by slightly expanding. This effect is the source of the lower densities in the cores of CDM halos in the presence of constant couplings. However, if the coupling is changing in time, there is an additional mechanism coming into play: the total potential energy of the system decreases due to the increase of the effective gravitational constant as a consequence of the growing scalar fifth-force. In our models, the effective gravitational constant grows by $\\sim $30 - 75 \\% during the whole cosmic evolution, which determines a corresponding decrease of the gravitational potential energy of collapsed systems. This decrease of total energy then determines a contraction of the halos.  Interstingly, we have shown that the two opposite behaviors found for the inner overdensity of nonlinear structures are determined by the background evolution of the coupled DE-CDM system: in the models that do not present a ``Growing $\\phi $MDE\" phase the friction term, which is the main driver of the expansion of halos in constant coupling models, is strongly suppressed and cannot counteract the contraction induced by the strong increase of the effective gravitational constant. On the contrary, the exponential coupling models, due to the presence of a ``Growing $\\phi $MDE\", have a still efficient friction term that balances and sometimes overcomes the effect of the potential energy decrease, thereby determining an overall expansion of the halos.  We have therefore shown how the nonlinear behavior of matter particles at small scales, in the context of coupled DE models with time dependent couplings, is directly influenced by the cosmological background evolution of the scalar field, and how the presence of a ``Growing $\\phi $MDE\" phase is essential to determine whether halos will contract or expand in these cosmologies.  These considerations apply also to halo concentrations, which are found to be higher with respect to $\\Lambda $CDM in the phenomenological models due to the absence of a ``Growing $\\phi $MDE\", and lower in the exponential coupling models.\\\\  In conclusion, we have performed a  complete numerical study of interacting DE cosmologies for a few general types of time dependence of the DE-CDM coupling, concerning background evolution, linear perturbations evolution, and nonlinear structure formation. We have presented the first high-resolution hydrodynamical N-body simulations of structure formation in the context of variable coupling models to date. In our analysis, we have found that differently from the constant coupling case, halo density profiles and halo concentrations do not evolve in the same direction with respect to $\\Lambda $CDM for all types of coupling evolution. In particular, depending on the type of background evolution determined by the coupling function, density profiles can be less overdense and correspondingly less concentrated than in $\\Lambda $CDM, or vice versa. Furthermore, the growth of structures at large scales is also affected in a significantly different way according to the different types of coupling evolution. Finally, we find that the decrease of the halo baryon fraction already found for constant coupling models can be significantly enhanced in variable coupling cosmologies. Some of these effects alleviate tensions between astrophysical observations and the $\\Lambda $CDM cosmology at small scales, and arise in cosmological models that contrarily to the constant coupling scenarios are not in stark conflict with present observational constraints on the background evolution of the Universe even in the presence of a significant coupling strength at low redshifts. Therefore, cosmological models with time dependent couplings in the dark sector might represent -- for some specific forms of coupling evolution -- a viable alternative to the standard $\\Lambda $CDM concordance model."
    },
    "1005/1005.2141_arXiv.txt": {
        "abstract": "We present detailed observations of the bright short-hard gamma-ray burst GRB~090510 made with the Gamma-ray Burst Monitor (GBM) and Large Area Telescope (LAT) on board the \\Fermi\\ observatory.  GRB~090510 is the first burst detected by the LAT that shows strong evidence for a deviation from a Band spectral fitting function during the prompt emission phase.  The time-integrated spectrum is fit by the sum of a Band function with $\\Epeak = 3.9\\pm 0.3$\\,MeV, which is the highest yet measured, and a hard power-law component with photon index $-1.62\\pm 0.03$ that dominates the emission below $\\approx$\\,20\\,keV and above $\\approx$\\,100\\,MeV.  The onset of the high-energy spectral component appears to be delayed by $\\sim$\\,0.1\\,s with respect to the onset of a component well fit with a single Band function. A faint GBM pulse and a LAT photon are detected 0.5\\,s before the main pulse. During the prompt phase, the LAT detected a photon with energy $30.5^{+5.8}_{-2.6}$ GeV, the highest ever measured from a short GRB. Observation of this photon sets a minimum bulk outflow Lorentz factor, $\\Gamma\\ga$\\,1200, using simple $\\gamma\\gamma$ opacity arguments for this GRB at redshift $z = 0.903$ and a variability time scale on the order of tens of ms for the $\\approx$\\,100\\,keV--few MeV flux. Stricter high confidence estimates imply $\\Gamma \\ga 1000$ and still require that the outflows powering short GRBs are at least as highly relativistic as those of long duration GRBs.  Implications of the temporal behavior and power-law shape of the additional component on synchrotron/synchrotron self-Compton (SSC), external-shock synchrotron, and hadronic models are considered. ",
        "introduction": "A difficulty in trying to understand gamma-ray bursts is that, at least in terms of the temporal structure of their emission, all GRBs differ.  When the overall time scales of the emission are considered, however, a pattern does emerge.  The durations of the $\\sim 100$\\,keV--MeV emission from GRBs form a bimodal distribution and hence are divided into two classes, namely the short and long duration bursts. The short bursts formally have durations $< 2\\,$s with typical values around $\\sim 0.2$\\,s, whereas the long bursts have a distribution that peaks around $\\sim 30$\\,s with a tail extending to several hundreds of seconds \\citep{Kouveliotou93}.  Although the physics of their gamma-ray emission is not well understood, these two classes of bursts likely originate from distinct types of progenitor systems \\citep{Lee07, Woosley06}. Long duration bursts are thought to be produced by the core collapse of massive stars as evidenced by the direct association of several nearby GRBs ($z<0.3$) with SN Ib/c events \\citep{Woosley06}.  Consistent with this, the afterglow counterparts of long duration GRBs tend to lie in star forming regions of low mass, irregular galaxies \\citep{Kocevski09}.  By contrast, short duration GRBs have been associated with both early and late type host galaxies, in proportions that reflect the underlying field galaxy distribution \\citep{Berger2009}. In the prevailing model for short bursts, they are produced in merger events of a compact binary systems composed of two neutron stars or a neutron star/black hole pair and so would tend to originate from older stellar populations.  With the launch and successful operation of the \\Fermi\\ Gamma-ray Space Telescope, a wider observational window has been opened through which a greater understanding of GRBs may be obtained.  The Large Area Telescope aboard \\Fermi\\ provides significantly greater energy coverage (20 MeV to $>$\\,300\\,GeV), field-of-view ($2.4$\\,sr) and effective area ($8000$\\,cm$^2$ at 1\\,GeV) than its predecessor EGRET \\citep{2009ApJ...697.1071A}.  Owing to its substantially lower deadtime (26\\,$\\mu$s vs 100\\,ms for EGRET), the LAT can probe the temporal structure of even the shortest GRBs.  In addition, the LAT can localize GRBs with sufficiently high precision to enable follow-up observations by \\Swift\\ and ground-based observatories; and at energies $\\ga$\\,few\\,GeV, the LAT can distinguish GRB photons from background with little ambiguity.  The Gamma-ray Burst Monitor, the other science instrument on \\Fermi, comprises an array of 12 NaI scintillators and two BGO scintillators and can detect gamma-rays from 8\\,keV to 40\\,MeV over the full unocculted sky. The combined capabilities of the LAT and GBM enable \\Fermi\\ to measure the spectral parameters of GRBs over seven decades in energy.  For sufficiently bright bursts, time-resolved spectral analysis is possible over the entire energy range.   The \\Fermi\\ observations of the short gamma-ray burst, GRB~090510, take full advantage of the GBM and LAT capabilities.  The LAT emission shows temporal structure on time scales as short as 20\\,ms. In addition to the usual Band function component \\citep{Band:93}, spectral fits reveal a hard power-law component emerging in the LAT band 0.1\\,s after the onset of the main prompt emission in the GBM band.  Moreover, a $\\approx 0.2$\\,s delay is observed between the brightening of the $\\approx 200$ MeV--GeV emission with respect to the strong count increases in the NaI and BGO. These behaviors present severe challenges for emission models of GRBs.  A photon with energy $30.5_{-2.6}^{+5.8}$\\,GeV was detected by the LAT 0.829\\,s after the GBM trigger.  This event arrived during the prompt phase and is temporally coincident with a sharp feature in the GBM and LAT light curves.  Given this energy, the temporal structure of the burst light curve and the known distance to the burst, \\citet{Nature..GRB090510} have used this photon and its arrival time to set limits on a possible linear energy dependence of the propagation speed of photons due to Lorentz-invariance violation that would require a quantum-gravity mass scale significantly above the Planck mass.  Similar to several of the long bursts seen by the LAT, GRB~090510 shows a high energy extended emission component that is detected by the LAT as late as 200\\,s after the GBM trigger. In the context of GRB outflow models, the properties of this GeV emission and the optical and X-ray afterglow observations by \\Swift\\ place significant constraints on possible internal and external shock models for the late time emission of this source \\citep{2009arXiv0909.0016G,2010ApJ...709L.146D,2009arXiv0911.4453C,2009ApJ...706L..33G,2009MNRAS.400L..75K}.  In this paper, we report on and discuss the GBM and LAT observations of GRB~090510 during the prompt emission phase.  In section~2, we give the basic observational details.  In section~3, we present a timing analysis, including discussion of its designation as a short-hard GRB and a cross-correlation analysis between various energy bands to characterize any energy-dependent temporal lags.  In section~4, we perform spectral analyses of both the time-integrated and the time-resolved data; and we demonstrate the significance of the additional power-law component and how the spectra evolves over the course of the burst.  In section~5, we derive a lower limit on the bulk Lorentz factor given the variability time scales and observations of the highest energy photons, describe the constraints imposed on leptonic, hadronic and other models of the prompt emission, and discuss implications of the detection of the 30.5\\,GeV photon for models of the extragalactic background light (EBL).  Finally, we summarize our results in section~6. ",
        "conclusions": "\\label{sec:discussion}  The emergence of a distinct high-energy spectral component in the prompt-phase spectrum of GRB~090510 establishes that a hard emission component in addition to a Band component is found in the short hard class of GRBs.  Hard power-law components are also found in three long-duration GRBs, namely GRB~090902B \\citep{Fermi_GRB090902B}, GRB~090926A \\citep{Fermi..GRB090926A}, and GRB~941017 \\citep{Gonzalez:03}.  In GRB~090510, the LAT data show that the $\\gamma$-ray flux above 100 MeV brightens $\\approx 0.2$\\,s after the start of the bright phase of GBM emission.  This behavior is similar, but on a shorter timescale, to the delayed onset in the long-duration GRB~080916C \\citep{2009Sci...323.1688A} and most other \\Fermi\\ GRBs, including GRB~081024B, the first short GRB observed with the LAT \\citep{2010ApJ...712..558A}.  Furthermore, GRB~090510 displays $>$\\,100\\,MeV emission significantly extended beyond the duration of the GBM flux \\citep{2010ApJ...709L.146D}, as observed in other \\Fermi\\ GRBs and earlier from GRB 940217 with the EGRET instrument on the {\\it Compton Observatory} \\citep{Hurley:94}.  These behaviors provide important constraints for high-energy emission models and could help answer whether the high-energy $\\gamma$ rays have a leptonic or hadronic origin.  Though the afterglow radiation in both long and short GRBs is probably nonthermal synchrotron emission from an external shock \\citep{1998ApJ...497L..17S}, the situation is less clear in the prompt and early afterglow phases when the GRB engine is most powerful.  This radiation could be from the thermal photosphere made by the powerful relativistic wind \\citep{2002ApJ...578..812M,2007ApJ...664L...1P}, from magnetic reconnection in Poynting-flux dominated outflows \\citep{2003astro.ph.12347L}, or from nonthermal leptonic emissions formed by internal or external shocks \\citep[e.g.,][]{2009MNRAS.400L..75K,2009arXiv0909.0016G,2009arXiv0911.4453C} in the relativistic jet of a GRB.   \\citet{2010ApJ...709L.146D} describe and present models for the \\Swift\\ and \\Fermi\\ observations of GRB~090510 during the afterglow phase.  Here we consider the implications of the prompt phase and early afterglow emission for GRB 090510.  After deriving the minimum bulk Lorentz factor $\\Gamma_\\min$ and considering the various uncertainties that enter into this calculation, we use the observations to constrain leptonic synchrotron/SSC model and hadronic models of short duration GRBs.  We do not discuss a thermal photospheric interpretation for the \\Fermi\\ results on GRB~090510. The photospheric interpretation overcomes the problem that the GBM spectra are harder than expected below $\\Epeak$ with the simplest synchrotron emission model (which is only the case at $\\gtrsim 2\\sigma$ in GRB~090510 during interval b; see Table~\\ref{tab:spectral fits}).  Even if it explains much of the GBM emission, however, a different origin is needed for the separate hard spectral component observed at LAT energies.  The coincident narrow spikes between the LAT all events and GBM lightcurves would not be easy to explain in a purely photospheric scenario, though Compton-scattered photospheric emission by internal shocked electrons could produce the coincident components \\citep{2010arXiv1002.2634T}.   \\subsection{Lower Limit on the Bulk Lorentz Factor}  The use of $\\gamma$-ray observations to constrain the bulk outflow speed of highly variable and energetic $\\gamma$-ray emission from GRBs has been studied by many authors \\citep[e.g.,][]{1997ApJ...491..663B,2001ApJ...555..540L,2004ApJ...613.1072R}. A detailed derivation involves an integration over the photon spectrum to calculate the opacity of $\\gamma$ rays emitted from sources with idealized geometries \\citep[see Supplementary Information in][]{2009Sci...323.1688A}.  A $\\delta$-function approximation for the $\\gamma\\gamma$ opacity constraint gives values of $\\Gamma_\\min$ accurate to $\\sim 10$\\% whenever the target photon spectrum is softer than $\\nu F_\\nu \\propto \\nu$.  In this case, $\\gamma\\gamma$ opacity arguments for a $\\gamma$-ray photon with energy $\\e_1$, in $m_ec^2$ units, imply a minimum bulk Lorentz factor, defined by $\\tau_{\\gamma\\gamma}(\\e_1) = 1$, of \\citep[e.g.][]{1995MNRAS.273..583D,2007PhR...442..166N}, \\begin{equation} \\Gamma_{\\rm min} \\cong \\left[ {\\sigma_{\\rm T} d_L^2 (1+z)^2 f_{\\hat \\e}\\e_1\\over 4t_v m_ec^4}\\right]^{1/6}\\;\\;,\\;\\; \\hat \\e = {2\\Gamma^2\\over (1+z)^2\\e_1}\\;. \\label{eqs2} \\end{equation} Here $f_\\e $ is the $\\nu F_\\nu$ flux at photon energy $m_ec^2 \\epsilon$, which is evaluated at $\\e = \\hat \\epsilon$ due to the peaking of the $\\gamma\\gamma$ cross section near threshold.  While the local value of the photon index around $\\hat\\epsilon$ has some effect on the exact numerical coefficient, this effect is small provided that the target photon index is $< -1/2$. Because of the threshold condition used to relate the high-energy photon and the target photons, the solution to equation~(\\ref{eqs2}) is iterative but quickly converges.  We use this expression to estimate $\\Gamma_\\min$ from \\Fermi\\ observations of GRB 090510 for comparison with more accurate calculations.  For interval b, during which a 3.4\\,GeV photon was detected, spectral analysis of GBM and LAT data during this episode reveals distinct Band-function and power-law components (Figure~\\ref{fig:nuFnu spectra}). The Band function has $\\Epeak = 5.1$ MeV, $\\alpha = -0.48$ and $\\beta = -3.09$ (Table~\\ref{tab:spectral fits}). The combined Band plus power-law fit reaches a peak $\\nu F_\\nu$ flux of $\\approx 4\\times 10^{-5}$ erg cm$^{-2}$ s$^{-1}$.  Writing the variability timescale $t_v$(s)$=0.01 t_{-2}$ s for 10 ms variability timescale, and $f_{\\hat \\e} = 10^{-5} f_{-5}$ erg cm$^{-2}$ s$^{-1}$, then equation~(\\ref{eqs2}) gives $\\Gamma_\\min \\cong 1100(f_{-5}/t_{-2})^{1/6}\\equiv 10^3\\Gamma_3$ with $\\epsilon_1 = 3400/0.511 \\cong 6650$. The target photon energy $\\hat \\e \\cong 2\\Gamma_\\min^2/(1+z)^2\\e_1 \\cong 110 (f_{-5}/t_{-2})^{1/3}$, or $\\approx 50$ MeV, corresponding to the Band $\\beta$ branch of the function.  Depending on whether the 3.4\\,GeV photon is interacting with the total emission or just the photons in the power law, then $f_{-5} \\approx 0.7$ or $f_{-5}\\approx 0.1$, and $\\Gamma_\\min \\cong 950$ or $\\Gamma_\\min \\cong 720$, respectively. For interval c from \\interval{0.8}{0.9}, the same procedure with the 30.5\\,GeV photon gives $\\Gamma_\\min \\cong 1370$ or 1060 for $f_{-5} \\approx 0.5$ or $f_{-5}\\approx 0.1$ corresponding respectively to the combined Band plus power-law fit or the power-law component only.  The results of numerical integrations to determine $\\Gamma_\\min$ using the more detailed expressions in \\citet{2009Sci...323.1688A} are shown in Table~\\ref{tab:Gamma_min_results}.  As can be seen, the simple estimates given above are in good agreement with the detailed calculation.  A number of issues arise in the use of equation~(\\ref{eqs2}) or the numerical integrations that are important for assessing the value and uncertainty in $\\Gamma_\\min$. These include the error incurred by the uncertainties in source spectral fitting parameters, which properly involves a covariance matrix to correlate uncertainties for different parameters of the Band-function and power-law fit.  For $E_{\\rm max}$, we take the highest energy photon associated with that pulse. Table~\\ref{tab:Gamma_min_results} presents the values for $t_v$, $E_{\\rm max}$, and $\\Gamma_\\min$ for time intervals \\interval{0.6}{0.8} and \\interval{0.8}{0.9}, which are the only two for which a distinct pulse width could be measured.  In interval \\interval{0.6}{0.8}, the Band + PL fit shown in Table~\\ref{tab:spectral fits} form the target photon spectrum, while in interval \\interval{0.8}{0.9}, we present results for both the Band and Band + PL fits since each fits the data reasonably well. We also give values for $\\Gamma_\\min$ assuming that only the hard power-law component forms the target photon source. Here we assume that $t_v$ is the same as that measured from the BGO emission, which is primarily associated with photons in the Band portion of the spectrum. If the variability timescale of the power-law emission is different than assumed, which would be compatible with the two components originating from different locations, then the minimum Doppler factor limit would change as indicated by equation~(\\ref{eqs2}). Furthermore, for calculations of $\\Gamma_\\min$ we use the spectrum derived on 0.2 s (time interval b) and 0.1 s (time interval c) timescales rather than on the shorter variability timescales during which the high-energy photons are measured. This is required for accurate spectral analysis, but could underestimate the flux (and therefore $\\Gamma_\\min$) during the bright narrow spikes, as can be seen from Fig. \\ref{fig:light curves}.  The derivation of $\\Gamma_\\min$ depends crucially on the assumption that the high-energy radiation and target photons are made in the same emitting region.  Correlated variability between different wavebands would support the cospatial assumption (see Figure~\\ref{fig:light curves}), but no strong evidence for this behavior was found in the \\ccf\\ analysis described in Section~3. A conservative assumption would be to suppose that the high-energy photon is part of the power-law component and that it can potentially interact only with target photons that are part of the same power-law emission component. Even in cases where the target MeV photons are made at smaller radii than the high-energy photons, or in different regions within the Doppler cone of the emitting surface, spacetime overlap will add to opacity, so this should represent the most conservative assumption. Further complicating the derivation of $\\Gamma_\\min$ is the assumed emitting geometry and the temporal evolution of the radiating plasma. For a blast-wave geometry, the precise value of $\\Gamma_\\min$ depends on whether high-energy photons are produced throughout the ``shell'' or from the inner edge of the ``shell,'' and on the dynamical behavior of the target photons \\citep{Granot:08}.  Finally, a significant uncertainty on $\\Gamma_\\min$ can arise if the photon with observed energy $E_{\\rm max}$ is a random fluctuation of the underlying true spectrum that corresponds to $\\Gamma \\lesssim \\Gamma_\\min$ and $\\tau_{\\gamma\\gamma}(E_{\\rm max}) \\gtrsim 1$. The confidence we have on the value of $\\Gamma$ depends on the radiative transport and escape of $\\gamma$ rays from the emitting region. For interval c from \\interval{0.8}{0.9}, $\\Gamma_\\min = 1218 \\pm 61$ for the Band plus power-law fit (Table~\\ref{tab:Gamma_min_results}). Assuming that the intrinsic spectrum extrapolates as a power law to high-energies, a likelihood ratio test assuming an exponential escape probability gives $\\Gamma/\\Gamma_\\min = 0.96$, $0.88$, and $0.80$ and a spherical escape probability gives $\\Gamma/\\Gamma_\\min = 0.89$, $0.69$, and $0.49$ at the 1, 2, and $3\\sigma$ confidence levels, respectively. The presence of two photons with energies between 1 and 2 GeV in interval b and a 7 GeV photon in interval c reduces the likelihood that the highest energy photon is a fluctuation and can be used to independently estimate $\\Gamma_\\min$, giving a value $\\Gamma_\\min \\gtrsim 1000$.  GRB 090510 is the second short GRB observed with LAT, after GRB~081024B \\citep{2010ApJ...712..558A}, but the first with a redshift, which is required to derive $\\Gamma_\\min$. The value of $\\Gamma_\\min \\cong 1200$--1300 for GRB 090510 is comparable to, and slightly larger than the values of $\\Gamma_\\min \\cong 900$ and $\\Gamma_\\min \\cong 1000$ derived for GRB 080916C \\citep{2009Sci...323.1688A} and GRB 090902B \\citep{Fermi_GRB090902B} using corresponding $\\gamma\\gamma$ opacity arguments.  This has led to suggestions that the GRBs with the most luminous LAT emission are those with the largest bulk Lorentz factors \\citep{2009MNRAS.400L..75K,2009arXiv0909.0016G}.   \\subsection{Models for the Prompt Radiation from GRB 090510}  In addition to the requirement of bulk outflow Lorentz factors $\\Gamma \\gtrsim \\Gamma_\\min$, models for GRB 090510 should explain the $\\approx 0.2$\\,s delay of the onset of the $\\gtrsim 100$\\,MeV emission compared to the start of the main GBM emission at \\Time{0.5}, the appearance of a hard component, and the high-energy radiation extending to $\\approx$\\,\\Time{150}.  \\subsubsection{Synchrotron/SSC Model}  A standard GRB model for the prompt phase assumes that the keV--MeV emission is nonthermal synchrotron radiation from shock-accelerated electrons \\citep[e.g.,][]{1996PhRvL..76.3478T}.  This emission is necessarily accompanied by SSC radiation.  The SSC component is stronger for a large ratio of nonthermal electron to magnetic-field energy density, as expressed by the condition $\\epsilon_e \\gg \\epsilon_B$ \\citep{2001ApJ...548..787S,Zhang:01}, and  also when the GRB has a lower bulk Lorentz factor for an external shock origin. Here $\\epsilon_e$ and $ \\epsilon_B$ are the fractions of shocked energy transferred to nonthermal lepton and magnetic-field energy, respectively.  Lower bulk Lorentz factors that give stronger SSC components result in greater attenuation of high-energy $\\gamma$ rays from $\\gamma\\gamma$ pair-production processes, as discussed in Section 5.1.  For generic ($E_{\\rm iso} \\sim 10^{52}$ erg, $\\Gamma_0\\sim 300, n\\sim 1$ cm$^{-3}$) parameters expected in an external shock model, a hardening of the LAT spectrum due to the deceleration of the blast wave and the emergence of the SSC component in the LAT band was expected to take place in the afterglow phase, but not in the prompt phase \\citep{Dermer:00b}.  With the larger initial Lorentz factors $\\Gamma_0\\gtrsim 10^3$ implied by $\\gamma\\gamma$ arguments for LAT GRBs and the earlier emergence of an external shock component, an SSC component would persist after the decline of the synchrotron component due to less scattering taking place in the Klein-Nishina regime as the blast wave decelerates, and to the peak frequency of the SSC flux decreasing from the TeV to the GeV range. The detailed observations of GRB~090510 can help determine whether the hard power-law component appearing at $\\approx$\\,\\Time{0.7} can be explained by SSC emission in the LAT waveband during the prompt phase.  Figure~\\ref{fig:dermerSSC} shows results for a numerical model where synchrotron peak energy, peak flux and variability time are made to correspond to the observed values shown in Figure~\\ref{fig:nuFnu spectra}. This code employs a Compton kernel that accurately treats Compton-scattering of relativistic electrons throughout the Thomson and Klein-Nishina regimes, internal $\\gamma\\gamma$ opacity, synchrotron self-absorption, radiative escape for $\\gamma$ rays described by exponential and spherical escape probabilities, and second-order SSC.  The code does not include reprocessing of internally-absorbed radiation or effects of attenuation of $\\gamma$-ray photons by the EBL (see Section 5.3).  Parameters appropriate to the Band component in interval b (see Figures~\\ref{fig:light curves} and~\\ref{fig:nuFnu spectra}) are used, as shown by the dotted curves.  The parameters are $t_v = 14$\\,ms, $\\Epeak = 5.10$\\,MeV, $E_\\max = 3.4$\\,GeV, $\\alpha = -0.48$, $\\beta = -3.09$, and peak synchrotron flux $\\approx 4\\times 10^{-5}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ (see Figure~\\ref{fig:nuFnu spectra} and Tables~\\ref{tab:spectral fits} and~\\ref{tab:Gamma_min_results}). Figure~\\ref{fig:dermerSSC} shows results for $\\Gamma = 500$ (dashed curves) and $\\Gamma = 1000$ (solid curves), and for magnetic fields $B^\\prime = 1$\\,kG and $B^\\prime = 1$\\,MG in the upper and lower panels, respectively.  Note that subtraction of an underlying hard component will not change the peak $\\nu F_\\nu$ flux value by more than $\\approx 5$\\%.  The unattenuated power of the SSC component is comparable to the synchrotron power when $B^\\prime = 1$ kG, whereas the SSC component is much weaker in the strong-field case. The isotropic jet power for the $B^\\prime = 1$ kG model is $\\cong 2.0\\times 10^{54}$ erg s$^{-1}$ and $\\cong 5.8\\times 10^{53}$ erg s$^{-1}$ for $\\Gamma = 500$ and $\\Gamma = 1000$, respectively, and is dominated by the energy in the escaping radiation. When $B^\\prime = 1$ MG, the isotropic jet power is $\\cong 1.1\\times 10^{55}$ erg s$^{-1}$ and $\\cong 7\\times 10^{56}$ erg s$^{-1}$ for $\\Gamma = 500$ and $\\Gamma = 1000$, respectively, and is dominated by magnetic-field energy.  For $B^\\prime \\ll$ 1 kG, the SSC flux becomes much brighter than the synchrotron flux, and the jet power becomes dominated by particles, even assuming that all particle energy is in the form of relativistic electrons.  Because the SSC component is strongly attenuated by $\\gamma\\gamma$ processes for the $\\Gamma$ factors considered, an electromagnetic cascade will be formed with $\\gamma$ rays emerging at lower energies where the system becomes optically thin.  The timescale for the electromagnetic radiation to cascade to energy $E_\\gamma$ is shorter than the synchrotron time, given by $t_{\\rm syn} \\approx 0.006 (\\Gamma/1000)^{1/2}/[(B^\\prime/{\\rm kG})^{3/2} \\sqrt{E_\\gamma/100{\\rm~MeV}}]$ s.  Unless $B^\\prime\\ll 1$ kG, in which case much more power is found in the cascading SSC emission than in the synchrotron emission, the cascading timescale is too short to explain the $\\approx 0.2$ s delay between the GBM and LAT emission.  This model also faces the well-known line-of-death problem \\citep{Preece:98} that the standard synchrotron mechanism makes a spectrum softer than $\\alpha = -2/3$, whereas $\\alpha = -0.48\\pm 0.07$ in interval b, representing a nearly $3\\sigma$ discrepancy from the hardest expected synchrotron emissivity. For the strong-field case where the SSC component is weak, the separate hard component in GRB 090510 would then have to originate from a different mechanism.  \\subsubsection{Afterglow Synchrotron Model}  \\citet{2009MNRAS.400L..75K,2009arXiv0910.5726K} and \\citet{2009arXiv0909.0016G,2009arXiv0910.2459G} have proposed forward shock emission from the early afterglow as the origin of the delayed onset and the hard component extending into the LAT energy band. This possibility is also considered by \\citet{2010ApJ...709L.146D}, \\citet{2009arXiv0911.4453C}, and \\citet{2009ApJ...706L..33G}. In particular, \\citet{2009arXiv0909.0016G} calculate the coasting bulk Lorentz factor of the GRB jet by identifying the time of the peak LAT emission occurring at $\\approx$\\,\\Time{0.7} with the deceleration time of a relativistic blast wave.  If $\\eta$ is the efficiency to convert bulk kinetic energy to $\\gamma$-ray energy, $\\Gamma_0 \\approx 2000 ~n^{-1/8} (t_{\\rm peak}/0.2~\\rm s)^{-3/8} (\\eta/0.2)^{-1/8} (E_{\\gamma,\\rm iso}/3.5\\times 10^{52}~\\rm erg)^{1/8}$, where $n($cm$^{-3}$) is the density of the surrounding medium. This expression uses an apparent isotropic $\\gamma$-ray energy release that excludes the LAT emission, which if due to synchrotron radiation from relativistic electrons, would require a radiative efficiency approaching unity.  Depending on circumburst density, the implied Lorentz factor is about 2--4 times the value of $\\Gamma_\\min$ calculated in the previous section. For this model, the emission radius corresponding to the time of the peak LAT flux is $R \\approx 2.4\\times 10^{16} n^{-1/4}$ cm.  At $t\\approx t_{\\rm peak}$, the minimum energy electrons in the forward shock radiate synchrotron photons at energies \\\\ $h\\nu_m \\approx 3.6(\\epsilon_e/0.1)^2\\xi_e^{-2}(\\epsilon_B/0.01)^{1/2}E_{53}^{1/2}(t/0.2\\,{\\rm s})^{-3/2}\\;$MeV assuming an electron injection index of $p \\approx 2.4$, where $\\xi_e$ is the fraction of the electrons that take part in the non-thermal power-law component responsible for the observed emission.  Recent particle in cell simulations of relativistic collisionless shocks \\citep{2008ApJ...673L..39S,2008ApJ...682L...5S,2009ApJ...695L.189M} suggest that $\\xi_e$ is fairly small (of the order of a few percent), which may in our case allow $h\\nu_m$ of several MeV even for low values of $\\epsilon_B$. For $\\epsilon_e = 0.1$, $\\epsilon_B = 0.01$, and an isotropic equivalent kinetic energy of $E_{k,{\\rm iso}} = 10^{53}\\;$erg (similar to $E_{\\rm\\gamma,iso}$), the cooling break frequency is given by \\citep{2002ApJ...568..820G}, $h\\nu_c \\approx 0.4 n^{-1} (\\epsilon_e/0.1)^{-1} (\\epsilon_B/0.01)^{-1/2} E_{53}^{-1/2} (t/0.2{\\rm~s})^{-1/2}$\\,keV. This expression holds when $\\epsilon_B \\ll \\epsilon_e$ and Klein-Nishina effects are unimportant, so that $Y = [(1+4\\epsilon_e/\\epsilon_B)^{1/2}-1]/2 \\approx \\sqrt{\\epsilon_e/\\epsilon_B} > 1$. This would imply fast cooling $\\nu_c \\ll \\nu_m$ for $n \\sim 1$, which could result in a highly radiative shock for $\\epsilon_e \\sim 1$. However, the late time broad band spectrum at $t \\sim 100\\;$s from optical-UV through X-ray and up to the LAT $\\gamma$-ray energies \\citep{2010ApJ...709L.146D} suggests $h\\nu_c(100\\,{\\rm s}) \\sim 300\\;$MeV, which, even if with a large uncertainty, together with the overall afterglow modeling suggest a much lower external density of $n \\sim 10^{-5}$. Such a low density would imply $h\\nu_c \\approx 0.3 (n/10^{-5})^{-1} (\\epsilon_B/0.01)^{-3/2} E_{53}^{-1/2} (t/0.5{\\rm~s})^{-1/2}$\\,GeV (where the Klein-Nishina effect suppresses SSC cooling), so that $\\nu_c$ passes through the (low part of the) LAT energy range around $\\sim 0.5\\;$s from the onset of the main emission episode, or at $\\sim T_0+1\\;$s, when there is a softening in the LAT photon index, near the end of the prompt emission. Thus, after the prompt emission, the LAT energy range would be above both $\\nu_m$ and $\\nu_c$, accounting for the observed photon index. This would imply $h\\nu_c \\approx 18 (n/10^{-5})^{-1} (\\epsilon_B/0.01)^{-3/2} E_{53}^{-1/2} (t/100{\\rm s})^{-1/2}$\\,MeV which is consistent with the broad band spectrum at that time, especially since the spectral break around $\\nu_c$ is very smooth and gradual \\citep{2002ApJ...568..820G}. Interestingly, the inferred value of $h\\nu_m(100\\,{\\rm s}) \\approx 0.43\\;$keV and the $\\nu_m \\propto t^{-3/2}$ scaling gives $h\\nu_m(t_{\\rm peak}\\approx 0.2\\,{\\rm s}) \\approx 4.8\\;$MeV, which is very close to the measured value of $\\Epeak$ near $t_{\\rm peak}$. This model would not produce a spectrum softer than $\\nu F_{\\nu} \\propto \\nu^{4/3}$, so would have difficulty accounting for the emission near 10\\,keV in interval b, which appears to be a continuation of the hard spectral component.  For an adiabatic blast wave and an electron injection index $p$, the synchrotron flux scales as $\\nu F_\\nu\\propto t^{(2-3p)/4} \\nu^{(2-p)/2}$ at frequencies $\\nu > \\nu_m,\\nu_c$. The measured late time LAT flux decay rate of $\\nu F_\\nu \\propto t^{-1.38\\pm 0.07}\\nu^{-0.1\\pm 0.1}$ would in this case require $p = 2.5\\pm 0.1$, while the spectral slope requires $p = 2.2\\pm 0.2$. Both are consistent with $p = 2.4$ at the $1\\sigma$ level. This value is also consistent with the late time afterglow of GRB~090510 \\citep{2010ApJ...709L.146D}. A radiative blast wave at early times is not required in order to account for the observed early LAT flux decay rate.   \\subsubsection{Hadronic Models}  In hadronic models, photohadronic and proton/ion synchrotron processes induce electromagnetic cascades, which lead to synchrotron and Compton emissions from secondary electron-positron pairs \\citep[e.g.,][]{2006NJPh....8..122D, Gupta:07, 2007ApJ...671..645A}. For a target photon energy distribution $n(\\epsilon) \\propto \\epsilon^{x}$, the efficiency for photopion processes is $ \\propto R^{-1} \\Gamma^{-2} E_{\\rm p}^{-1-x} \\propto \\Gamma^{-4} t_{\\rm v}^{-1} E_{\\rm p}^{-1-x}$, where $R\\propto c\\Gamma^2 t_{\\rm v}$ is the shock radius. In this expression, protons with energy $E_{\\rm p}$ preferentially interact with photons with energies $\\propto \\Gamma^2/E_{\\rm p}$ \\citep{1995PhRvL..75..386W, 2009PhRvL.103h1102M}.  The large deduced values for $\\Gamma$ in GRB~090510 make the photopion efficiency low, so that a very large energy release is required if the LAT radiation from GRB 090510 is assumed to be from a photomeson-induced cascade.  A stronger magnetic field shortens the acceleration timescale, leading to a larger maximum particle energy $E_{\\rm max}$, thus enhancing the photopion production efficiency.  In such a strong magnetic field, however, the effective injection index of secondary pairs tends to be about $-2$ \\citep{1992MNRAS.258..657C}, which yields a flat spectrum in a $\\nu F_\\nu$ plot, while the power-law index of the extra component in GRB~090510 is $\\sim -1.6$.  In a weaker magnetic field, the Compton component from secondary pairs can harden the spectrum, though with reduced photopion production efficiency due to the smaller value of $E_{\\rm max}$. The slower cooling time of protons than electrons would produce an extended proton-induced emission feature \\citep{Bottcher:98}, though the SSC component would decay even more slowly \\citep{Zhang:01}. A recent numerical calculation by \\citet{2009ApJ...705L.191A} for GRB 090510 demonstrates that the proton injection isotropic-equivalent power is required to be larger than $10^{55}$\\,erg\\,s$^{-1}$ to explain the hard spectra of the extra component in GRB 090510, which is $\\approx 2$ orders of magnitude greater than the measured apparent isotropic $\\gamma$-ray luminosity.  An alternative hadronic scenario is a proton synchrotron model with a very strong magnetic field \\citep{2009arXiv0908.0513R}.  This model explains the delayed onset by proton synchrotron radiation in the prompt phase due to the time required to accelerate, accumulate, and cool the ultrarelativistic protons.  For an onset times $t_{\\rm onset} \\approx 0.2$~s after the start of the GBM main emission at $T_0+0.5$~s, the required magnetic field is $B^\\prime \\approx 7.4\\times 10^5 ~(\\Gamma/1000)^{-1/3} (t_{\\rm onset}/0.1~\\rm s)^{-2/3} (E_\\gamma/100~\\rm MeV)^{-1/3} $~G in the shocked fluid frame.  The corresponding total energy release, for a two-sided jet beaming factor $f_b \\approx 1.5\\times 10^{-4}(\\theta_j/1~{\\rm deg})^2$, is ${\\cal E}\\approx 1.3\\times 10^{54}(\\Gamma/1000)^{16/3} (t_{\\rm onset}/0.1~\\rm s)^{5/3} (E_\\gamma/100~\\rm MeV)^{-2/3} (\\theta_j/1~{\\rm deg})^2$~erg \\citep[see][for GRB 080916C]{2009ApJ...698L..98W}. Thus values of $\\Gamma \\lesssim 1000$ and narrow jet angles $\\approx 1^\\circ$ are required for a proton synchrotron scenario, and such strong beaming is not clearly found in the short hard class of GRBs \\citep{2007PhR...442..166N}. Proton-dominated GRB models for ultra-high energy cosmic rays therefore are plausible only with low values of $\\Gamma$ and narrow jet collimation to reduce the total energy.  \\subsection{Implications for the Extragalactic Background Light}  The EBL is dominated by direct starlight in the optical/ultraviolet and by stellar radiation that is reprocessed by dust in the infrared. The EBL is difficult to measure directly due to contamination by zodiacal and Galactic foreground light \\citep{2001ARA&A..39..249H}.  For sources at sufficiently high redshifts, $\\gamma\\gamma$ absorption of high-energy $\\gamma$-rays by EBL photons can provide a means of constraining models of the EBL. Figure~\\ref{fig:taugg} shows the absorption optical depth, $\\tau_{\\gamma\\gamma}$, for various models of the EBL as a function of $\\gamma$-ray energy at the redshift $z=0.903$ of GRB 090510.  We have included curves for the two models of \\citet{2006ApJ...648..774S} as well as the fiducial model of \\citet{2009arXiv0905.1144G}, the best fit model of \\citet{2004A&A...413..807K}, the model by \\citet{2008A&A...487..837F}, and ``Model C'' of \\citet{2009arXiv0905.1115F}.  The models of \\citet{2006ApJ...648..774S} border on optically thick at 30.5\\,GeV; all other models considered here give a transmission probability of $e^{-\\tau_{\\gamma\\gamma}}\\gtrsim 0.85$.  The baseline and fast evolution models of \\cite{2006ApJ...648..774S} give transmission probabilities of 0.37 and 0.30, respectively.  Although a higher energy ($\\approx 30.5$\\,GeV) photon was found from this burst than for GRB 080916C ($13$\\,GeV; \\cite{2009Sci...323.1688A}), that burst was more constraining for EBL models due to its higher redshift, $z = 4.35 \\pm 0.15$ \\citep{2009A&A...498...89G}. However, the EBL does evolve with redshift, and this GRB provides an independent constraint from a later time and at a closer distance than GRB 080916C. The low optical depth for the highest energy photons in GRB~090510 justify neglecting EBL effects in Figure~\\ref{fig:dermerSSC}."
    },
    "1005/1005.0144_arXiv.txt": {
        "abstract": "{The observed relation between the X-ray radiation from active galactic nuclei, originating in the corona, and the optical/UV radiation from the disk is usually  described by the anticorrelation between the UV to X-ray slope $\\alpha_{ox}$ and the UV luminosity. Many factors can affect this relation, including:  i) enhanced X-ray emission associated with the jets of radio-loud AGNs, ii) X-ray absorption associated with the UV broad absorption line (BAL) outflows, iii) other X-ray absorption not associated with BALs, iv) intrinsic X-ray weakness, v) UV and X-ray variability, and non-simultaneity of UV and X-ray observations. The separation of these effects provides information about the intrinsic $\\alpha_{ox}-L_{UV}$ relation and its dispersion, constraining models of disk-corona coupling.} {We use simultaneous UV/X-ray observations to remove the influence of non-simultaneous measurements from the $\\alpha_{ox}-L_{UV}$ relation.} {We extract simultaneous data from the second XMM-Newton serendipitous source catalogue (XMMSSC) and the XMM-Newton Optical Monitor Serendipitous UV Source Survey catalogue (XMMOMSUSS), and derive the single-epoch $\\alpha_{ox}$ indices. We use ensemble structure functions to analyse multi-epoch data.} {We confirm the anticorrelation of $\\alpha_{ox}$ with $L_{UV}$, and do not find any evidence of a dependence of $\\alpha_{ox}$ on $z$. The dispersion in our simultaneous data ($\\sigma\\sim 0.12$) is not significantly smaller than in previous non-simultaneous studies, suggesting that ``artificial $\\alpha_{ox}$ variability'' introduced by non-simultaneity is not the main cause of dispersion. ``Intrinsic $\\alpha_{ox}$ variability'' , i.e., the true variability of the X-ray to optical ratio, is instead important, and accounts for $\\sim 30\\%$ of the total variance, or more. ``Inter-source dispersion\", due to intrinsic differences in the average $\\alpha_{ox}$ values from source to source, is also important. The dispersion introduced by variability is mostly caused by the long timescale variations, which are expected to be driven by the optical variations.} {} ",
        "introduction": "The relationship between the X-ray and optical/UV luminosity of active galactic nuclei (AGNs) is usually described in terms of the index $\\alpha_{ox}=0.3838 \\log(L_X/L_{UV})$, i.e., the slope of a hypothetical power law between 2500 \\AA\\ and 2 keV rest-frame frequencies. The X-ray and UV monochromatic luminosities are correlated over 5 decades as $L_X \\propto L_{UV}^k$, with $k\\sim0.5-0.7$, and this provides an anticorrelation $\\alpha_{ox} = a \\log L_{UV}$ + const, with $-0.2\\la a\\la -0.1$ \\citep[e.g.,][]{avni86,vign03,stra05,stef06,just07,gibs08}. One of the main results of these analyses is that QSOs are universally X-ray luminous and that X-ray weak QSOs are very rare \\citep[e.g.,][]{avni86,gibs08}, but it is not yet known if the same is true for moderate luminosity AGNs. UV photons are generally believed to be radiated from the QSO accretion disk, while X-rays are supposed to originate in a hot coronal gas of unknown geometry and disk-covering fraction. The X-ray/UV ratio provides information about the balance between the accretion disk and the corona, which is not yet understood in detail. The $\\alpha_{ox}-L_{UV}$ anticorrelation implies that AGNs redistribute their energy in the UV and X-ray bands depending on the overall luminosity, with more luminous AGNs emitting fewer X-rays per unit UV luminosity than less luminous AGNs \\citep{stra05}. It has been proposed that the anticorrelation can be caused by the larger dispersion in the luminosities in the UV than the X-ray band for a population with intrinsically uniform $\\alpha_{ox}$ \\citep{la-f95,yuan98}; however, more recent analyses based on samples with wider luminosity ranges confirm the reality of the relationship \\citep{stra05}. \\citet{gibs08} stressed the quite large scatter in the X-ray brightness of individual sources about the average relation and investigated the possible causes of the dispersion. Part of this scatter, usually removed \\citep[e.g.,][]{stra05,stef06,just07,gibs08} is caused by radio-loud quasars, which are relatively X-ray bright because of the enhanced X-ray emission associated with their jets \\citep[e.g.,][]{worr87}, and to broad absorption line (BAL) quasars, which are relatively X-ray faint \\citep[e.g.,][]{bran00} due to X-ray absorption associated with the UV BAL outflows. Additional causes of deviations from the average $\\alpha_{ox}-L_{UV}$ relation include: i) X-ray absorption not associated with BALs, ii) intrinsic X-ray weakness, and iii) UV and X-ray variability, possibly in association with non-simultaneous UV and X-ray observations. In particular, \\citet{gibs08} estimate that variability may be responsible for 70\\%-100\\% of the $\\alpha_{ox}$ dispersion, and that a few percent ($<2$\\%) of all quasars are intrinsically X-ray weak by a factor of 10, compared to the average value at the same UV luminosity. A large fraction of intrinsically X-ray weak sources would suggest that coronae may frequently be absent or disrupted in QSOs. An extreme case is PHL 1811, which is X-ray weak by a factor $\\sim$70, studied in detail by \\citet{leig07}, who propose various scenarios, including disk/corona coupling by means of magnetic reconnections, Compton cooling of the corona by unusually soft optical/UV spectrum, and the photon trapping of X-ray photons and their advection to the black hole. The influence of variability on the $\\alpha_{ox}-L_{UV}$ relation can be divided into two different effects: i) non-simultaneity of X-ray and UV measurements, which we call ``artificial $\\alpha_{ox}$ variability'', and ii) true variability in the X-ray/UV ratio, which we refer to as ``intrinsic $\\alpha_{ox}$ variability''. It is beneficial to analyse simultaneously acquired X-ray and UV data to eliminate the effect of the artificial variability and search for the intrinsic X-ray/UV ratio and/or its variability. On a rest-frame timescale of a few years, the optical/UV variability of QSOs has been estimated to be $\\sim$30\\% \\citep[e.g.,][]{gial91,vand04}, while X-ray variability has been estimated to be $\\sim$40\\% for Seyfert 1 AGNs \\citep{mark03}. On intermediate timescales, the relation between X-ray and optical/UV variability may be due to either: i) the reprocessing of X-rays into thermal optical emission, by means of irradiation and heating of the accretion disk, or ii) Compton up-scattering, in the hot corona, of optical photons emitted by the disk. In the former case, variations in the X-ray flux would lead optical/UV ones, and vice versa in the latter case. Cross-correlation analyses of X-ray and optical/UV light curves allow us to constrain models for the origin of variability. The main results obtained so far, on the basis of simultaneous X-ray and optical observations, indicate a cross-correlation between X-ray and UV/optical variation on the timescale of days, and in some cases delays of the UV ranging from 0.5 to 2 days have been measured \\citep{smit07}. Simultaneous X-ray/UV data can be obtained by the XMM-Newton satellite, which carries the co-aligned Optical Monitor (OM). The second XMM-Newton serendipitous source catalogue (XMMSSC) \\citep{wats09} is available online in the updated incremental version 2XMMi\\footnote{http://heasarc.gsfc.nasa.gov/W3Browse/xmm-newton/xmmssc.html}. The XMM-Newton Optical Monitor Serendipitous UV Source Survey catalogue (XMMOMSUSS) is also available online\\footnote{http://heasarc.gsfc.nasa.gov/W3Browse/xmm-newton/xmmomsuss.html}. We look for simultaneous measurements of the $\\alpha_{ox}$ index from XMM/OM catalogues, to provide at least partial answers to the following questions: how large is the effect of non-simultaneous X-ray/UV observations on the dispersion about the average $\\alpha_{ox}-L_{UV}$ relationship? Is there any spectral X-UV variability for individual objects? Do their $\\alpha_{ox}$ harden in the bright phases or vice versa? Which constraints do these measurements place on the relationship between the accretion disk and the corona?  The paper is organised as follows. Section 2 describes the data extracted from the archival catalogues. Section 3 describes the SEDs of the sources and the evaluation of the specific UV and X-ray luminosities. Section 4 discusses the $\\alpha_{ox}-L_{UV}$ anticorrelation and its dispersion. In Sect. 5, we present the multi-epoch data and discuss the intrinsic X/UV variability of individual sources. Section 6 provides notes about individual peculiar sources. Section 7 discusses and summarises the results.  Throughout the paper, we adopt the cosmology H$_{o}$=70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m}$=0.3, and $\\Omega_{\\Lambda}$=0.7. ",
        "conclusions": "The behaviour of $\\alpha_{ox}$, i.e., its dependence on luminosity and redshift, its dispersion and variability, are to be considered as symptoms of the relation between disk and corona emissions and their variabilities.  It is generally believed that variable X-ray irradiation can drive optical variations by means of variable heating of the internal parts of the disk on relatively short timescales, days to weeks, while intrinsic disk instabilities in the outer parts of the disk dominate on longer timescales, months to years, propagating inwards and modulating X-ray variations in terms of Compton up-scattering in the corona \\citep{czer04,arev06,arev09,papa08,mcha10}.  The structure functions of the light curves increase on long timescales both in the optical \\citep[e.g.,][]{dicl96,vand04,baue09} and X-rays \\citep[e.g.,][Vagnetti et al., in prep.]{fior98}. This, however, does not imply that the $\\alpha_{ox}$ SF also increases with time lag. Larger changes (on long time lags) in both X-ray and UV fluxes may occur without changes in the spectral shape (i.e., with constant $\\alpha_{ox}$). Our results shown in Fig. 9 indicate that this is not the case, i.e., that slope changes are indeed larger on longer timescales.  Moreover, it is evident from Fig. 9 that most of the dispersion about the $\\alpha_{ox}-L_{UV}$ relation is due, in the present sample, to variations on timescales from months to years, which are associated with optically driven variations, according to the general belief.  The $\\alpha_{ox}$ structure function does not distinguish between the hardening or softening of the optical to X-ray spectrum during brightening. This is instead described by the spectral variability parameter $\\beta=\\partial\\alpha/\\partial\\log F$ \\citep{trev01,trev02}, which can be adapted to the optical-X-ray case to become  \\begin{equation} \\beta_{ox}={\\delta\\alpha_{ox}\\over\\delta\\log L_{UV}}\\,, \\end{equation}  \\noindent where $\\beta_{ox}$ is the slope of the correlated variations $\\delta\\alpha_{ox}$ and $\\delta L_{UV}$, and describes whether a source hardens when it brightens or vice versa, i.e., if the X-ray luminosity increases more than the optical or less. For example, single source variations parallel to the $\\alpha_{ox}$ anticorrelation have a negative $\\beta_{ox}$, while variations perpendicular to the correlation have $\\beta_{ox}>0$. Both of these behaviours can be seen in Fig. 7.  Of course, the different behaviour of the sources in the $\\alpha_{ox}-L_{UV}$ plane may correspond to a different time sampling. Constraining physical models of the primary variability source and disk-corona coupling would require the analysis of $\\beta_{ox}$ as a function of the time lag. This analysis does not look feasible, i.e., to have statistical reliability, with the present sparse sampling. We can propose more conventional scenarios and note that since most of the variability, in the present sample, occurs on long ($\\sim 1$ year ) timescale, it is presumably associated with optically driven variations. Considering all the measured variations $\\delta\\alpha_{ox}$ and $\\delta L_{UV}$, we obtain the ``ensemble\" average $\\langle\\beta_{ox}\\rangle=-0.240$. The negative sign implies that, on average, a spectral steepening occurs in the brighter phase. This is, in fact, consistent with larger variations in the UV band, driving the X-ray variability. The value of $\\langle\\beta_{ox}\\rangle$ can be compared with the average slope of the $\\alpha_{ox}-L_{UV}$ relation, Eq. (13), indicating that the UV excess in the brighter phase (steepening) is larger than the average UV excess in bright objects respect to faint ones.  Finally we emphasise that, despite its limitations, the present analysis illustrates the feasibility of an ensemble analysis of the $\\alpha_{ox}-L_{UV}$ correlation, e.g., by considering the $\\beta$ parameter as a function of time lag. What is presently missing is an adequate simultaneous X-ray-UV sampling, at relatively short time lags,  of a statistical AGN sample. An ensemble analysis may provide important constraints even when the total number of observations does not allow us to carry out  a cross-correlation analysis of X-ray and UV variations of individual sources.  We summarise our main results as follows: \\begin{itemize} \\item we have studied the $\\alpha_{ox}-L_{UV}$ anticorrelation with simultaneous data extracted from the XMM-Newton Serendipitous Source catalogues; \\item we confirm the anticorrelation, with a slope (-0.178) slightly steeper w.r.t. \\citet{just07}; \\item we do not find evidence for a dependence of $\\alpha_{ox}$ on redshift, in agreement with previous authors \\citep[e.g.][]{avni86,stra05,stef06,just07}; \\item there appears to be a flatter slope to the anticorrelation at low luminosities and low redshifts, in agreement with previous results by \\citet{stef06}; \\item the dispersion in our simultaneous data ($\\sigma\\sim 0.12$) is not significantly smaller w.r.t. previous non-simultaneous studies  \\citep{stra05,just07,gibs08}, indicating that ``artificial $\\alpha_{ox}$ variability'' introduced by non-simultaneity is not the main cause of dispersion; \\item ``intrinsic $\\alpha_{ox}$ variability'' , i.e., true variability in the X-ray to optical ratio, is important, and accounts for $\\sim 30\\%$ of the total variance, or more; \\item ``inter-source dispersion\", due to intrinsic differences in the average $\\alpha_{ox}$ values from source to source, is also important; \\item the dispersion introduced by variability is mostly caused by the long timescale variations, which are expected to be dominated by the optical variations; the average spectral softening observed in the bright phase is consistent with this view; \\item distinguishing the trends produced by optical or X-ray variations may be achievable using the ensemble analysis of the spectral variability parameter $\\beta_{ox}$ as a function of time lag; crucial information would be provided by wide field simultaneous UV and X-ray observations with relatively short (days-weeks) time lags. \\end{itemize}"
    },
    "1005/1005.2363_arXiv.txt": {
        "abstract": "We present manganese abundances in 10 red-giant members of the globular cluster $\\omega$ Centauri; 8 stars are from the most metal-poor population (RGB MP and RGB MInt1) while two targets are members of the more metal rich groups (RGB MInt2 and MInt3). This is the first time Mn abundances have been studied in this peculiar stellar system. The LTE values of [Mn/Fe] in $\\omega$ Cen overlap those of Milky Way stars in the metal poor $\\omega$ Cen populations ([Fe/H]$\\sim$-1.5 to -1.8), however unlike what is observed in Milky Way halo and disk stars, [Mn/Fe] declines in the two more metal-rich RGB MInt2 and MInt3 targets. Non-LTE calculations were carried out in order to derive corrections to the LTE Mn abundances.  The non-LTE results for $\\omega$ Cen in comparison with the non-LTE [Mn/Fe] versus [Fe/H] trend obtained for the Milky Way confirm and strengthen the conclusion that the manganese behavior in $\\omega$ Cen is distinct. These results suggest that low-metallicity supernovae (with metallicities $\\le$ -2) of either Type II or Type Ia dominated the enrichment of the more metal-rich stars in $\\omega$ Cen. The dominance of low-metallicity stars in the chemical evolution of $\\omega$ Cen has been noted previously in the s-process elements where enrichment from metal-poor AGB stars is indicated. In addition, copper, which also has metallicity dependent yields, exhibits lower values of [Cu/Fe] in the RGB MInt2 and MInt3 $\\omega$ Cen populations. ",
        "introduction": "The abundance ratios of chemical elements which have different nucleosynthetic origins can be used to constrain the history of chemical evolution in different types of stellar populations. This is because stellar sources, such as massive stars that undergo core-collapse and explode as supernovae of Types II or Ib/Ic, or binary systems that become Type Ia supernovae (SNe Ia), or asymptotic giant branch (AGB) stars, produce and return to the interstellar medium different element abundance ratios on differing timescales.   Abundance ratios of certain elements can thus be used to explore which stars and in which proportions have contributed to chemical evolution within stellar populations in galaxies.  Of particular interest to chemical evolution of a stellar system are those elements whose yields may depend on the metallicity of the progenitor star.  Metallicity-dependent abundance patterns can retain `memory' of the metallicity distributions of their stellar populations, regardless of how the overall metallicity of the parent galaxy evolves.   One important element whose production may be metallicity dependent is manganese (Z=25), which falls within the iron peak.  Manganese is produced in both core collapse supernovae (CC SNe) and SNe Ia, but the relative amounts from these two sources are not well constrained. Manganese has one stable isotope, $^{55}$Mn;  this element is synthesized as a result of explosive incomplete silicon burning (see, e.g., Thielemann et al. 2007 for a more complete discussion).  Woosley \\& Weaver (1995), for example, model nucleosynthesis of massive stars exploding as supernovae of Type II (SNe II) and find increasing Mn yields with increasing metallicity. Convolving the Mn yields from the Woosley \\& Weaver (1995) models with a Salpeter mass function leads to values of [Mn/Fe] that decline steadily to -0.3 at [Fe/H]=-1 and then to -0.5 at [Fe/H]=-2. A decrease of [Mn/Fe] ratios with metallicity is reported in the LTE Mn I abundance analyses of Galactic stars such as Gratton (1989), Reddy et al. (2003; 2006), Johnson (2002), Cayrel et al. (2004), or Fetzing et al. (2007). Recent results from Bergemann \\& Gehren (2008), however, indicate that the LTE approximation underestimates the abundances based on Mn I lines by 0.1 - 0.4 dex, with non-LTE effects being most pronounced in stellar atmospheres with low metal content. In the metallicity range investigated by Bergemann \\& Gehren (2008), -2.5 $<$ [Fe/H] $<$ 0, the non-LTE [Mn/Fe] ratios in late-type stars are approximately solar. The flat [Mn/Fe] trend with [Fe/H] is in general agreement with the SNe II yields of Chieffi \\& Limongi  (2004) who do not predict a strong depletion of Mn relative to Fe at low metallicities.  Manganese yields from SNe Ia models can span a range of values depending on the explosion mechanism, such as slow or fast deflagrations, or deflagration-detonation events.  Iwamoto et al. (1999) investigate a variety of models and find no large differences in Mn/Fe yields (i.e., less than a factor of $\\sim$2) from models driven by both slow and fast deflagrations, as well as deflagration-detonations, having different initial metallicities. Badenes et al. (2008) computed nucleosynthesis from 4 delayed detonation models, with different metallicities, as well as one deflagration model.  This group finds that Mn yields decline with decreasing metallicity, despite the differences in the explosion mechanisms and initial conditions of the models.  In addition to the results for the Galactic thin disk, thick disk and halo noted above, McWilliam et al. (2003) added two other populations to the studies of manganese by measuring LTE [Mn/Fe] abundances for stellar members of the Sagittarius dwarf spheroidal galaxy and stars from the Galactic bulge.  Both systems exhibited somewhat different behaviors of [Mn/Fe] with [Fe/H], with the Sgr dwarf galaxy stars having values of [Mn/Fe] that fall below the general LTE trend found for Milky Way disk or halo stars at a given [Fe/H].   Bulge stars exhibit opposite behavior, with their values of [Mn/Fe] falling above the trend defined by the Milky Way disk and halo stars.  The goal of this study is to add another distinct stellar population to analyses of manganese and its chemical evolution in different Galactic environments. $\\omega$ Cen exhibits some peculiar characteristics in the nature of its chemical evolution, with perhaps the most striking being the large increase in the abundances of the heavy s-process elements (such as Ba or La) as the overall metallicity of cluster stars, measured by such elements as Fe, Ca, or Ti, increases (e.g. Norris \\& Da Costa 1995). $\\omega$ Centauri, although historically classified as a globular cluster, is now thought possibly to be a surviving remnant of a captured small galaxy, with multiple populations spanning a large range in metallicity (for a more detailed discussion see the review by Smith 2004). Recently,  Carretta et al. (2010) have pointed out similarities between the $\\omega$ Cen populations and those from the Sagittarius dwarf galaxy. Five distinct stellar populations each with a different metallicity have been identified by Pancino et al. (2000) and Sollima et al. (2005). These studies label the distinct red giant branches from the $\\omega$ Cen populations as RGB MP (metal-poor); RGB MInt1; RGB MInt2; RGB MInt3 (intermediate metallicities); and RGB-a (anomolous, with the highest metallicity).  Manganese abundances are presented here for the first time in $\\omega$ Cen stars, with the sample consisting of 10 targets; 8 red giants are from the most metal poor populations (RGB MP and RGB MInt1) and 2 stars are from the more metal rich RGB MInt2 and MInt3.  These $\\omega$ Cen stars have been analyzed in previous studies (Smith et al. 2000; Cunha et al. 2002), however Mn was not included in the analysis. ",
        "conclusions": "Manganese abundances have been measured for the first time in the peculiar globular cluster $\\omega$ Cen, with the analysis of 10 red giants spanning a range in metallicity from [Fe/H]= -1.9 to -0.9.  The analysis is based on Mn I lines using both LTE and non-LTE calculations to derive abundances. In addition, the possible effects of enhanced He abundances on derived Mn abundances were investigated for the more metal-rich $\\omega$ Cen giants and were found to be negligible.  The novel result is that two members from the more metal-rich populations with $\\omega$ Cen (RGB MInt2 and MInt3) exhibit low ratios of Mn/Fe in comparison to Galactic field stars, as well as other globular cluster stars at the same metallcitiy ([Fe/H]$\\sim$-1).  Differences between $\\omega$ Cen and the other Milky Way populations exist whether the comparison is made using LTE or non-LTE abundances.  The low abundances of Mn may indicate that low-metallicity progenitors to supernovae (of either core collapse or SNe Ia) dominated the production of manganese within $\\omega$ Cen.  This result for Mn is similar to what has been noted previously for the behavior of copper (which in some nucleosynthesis processes has metallicity-dependent yields) in $\\omega$ Cen (Cunha et al. 2002).  The behavior of Mn in the more extreme metal-rich $\\omega$ Cen population (RGB-a) remains to be probed.  In future studies it would be of interest to determine Mn abundances in the most metal-rich $\\omega$ Cen population; for example, by analyzing Mn in the 3 most metal-rich red giants studied to date (ROA 300, WFI22068, and WFI222679) by Pancino et al. (2002).  The behavior of the manganese abundances in these more metal-rich stars will provide further insight into both the origins of Mn and the nature of star formation within $\\omega$ Cen during the final throes of its chemical evolution.   MB thanks Dr. Frank Grupp for providing MAFAGS-OS model atmospheres for selected stars. This research was supported in part by the National Science Foundation (AST 06-46790 to KC and VVS).  DLL thanks the Robert A. Welch Foundation for support via grant F-634."
    },
    "1005/1005.2964_arXiv.txt": {
        "abstract": "The survival of unbound density substructure against orbital mixing imposes strong constraints on the slope of the underlying gravitational potential and provides a new test on modified gravities.  Here we investigate whether the interpretation that the stellar clump in Ursa Minor (UMi) dwarf spheroidal galaxy is a `dynamical fossil' is consistent with Modified Newtonian dynamics (MOND). For UMi mass models inferred by fitting the velocity dispersion profile, the stellar clump around the second peak of UMi is erased very rapidly, within $1.25$ Gyr ($6.5$ orbits), even with the inclusion of self-gravity. We find that the clump can hardly survive for more than $2$ Gyr even under more generous conditions. Alternative scenarios which could lead to a kinematically cold clump are discussed but, so far, none of them were found to be fully satisfactory. Our conclusion is that the cold clump in UMi poses a challenge for both $\\Lambda$CDM and MOND. ",
        "introduction": "The standard concordance cosmological model with cold dark matter ($\\Lambda$CDM model) is remarkably successful on scales larger than $1$ Mpc, but it faces challenges on smaller scales. For instance, it seems that the theory predicts a too cuspy density profile for the dark matter at the centres of galaxies (e.g., Trachternach et al.~2008). The MOdified Newtonian Dynamics (MOND) proposed by Milgrom (1983) has proven to be successful in reproducing the kinematics of spiral galaxies without any assumption of unseen matter (see Sanders \\& McGaugh 2002, for a review), from extremely low mass galaxies of low surface brightness (Milgrom \\& Sanders 2006) to high luminosity galaxies (Sanders \\& Noordermeer 2007). Gentile et al.~(2007) found that the observed rotation curves in tidal dwarf galaxies are quite naturally explained without any free parameters within MOND, and are inconsistent with the current $\\Lambda$CDM theory (see also Kroupa et al.~2005). If MOND is able to naturally account for all of the discrepancies faced by $\\Lambda$CDM on small scales, this would lend strong support to MOND.      In the Newtonian dark matter scenario, dwarf spheroidal galaxies (dSph's) require the largest mass-to-light ratios. Hence, dSph's provide a unique testing ground for the nature of dark matter and its alternatives (Gerhard \\& Spergel 1992; Milgrom 1995; Lokas et al.~2006; S\\'anchez-Salcedo et al.~2006; Angus 2008). However, the mass-to-light ratios inferred in MOND are very sensitive to uncertainties on the structural parameters, luminosities, distances and internal velocity dispersions. It is, therefore, important to explore other gravitational effects, which, in principle, may offer independent tests.  In the dark matter paradigm, the velocity dispersion profiles of the brightest dSphs are consistent with both a cuspy NFW halo and a cored dark halo. However, there exists some indirect evidence that dSph galaxies may possess a core (Kleyna et al.~2003, hereafter K03; Goerdt et al.~2006; S\\'anchez-Salcedo et al.~2006). In particular, K03 considered the survival of the cold density substructure detected in photometric data in Ursa Minor (UMi). This substructure appears as an off-centre localized stellar clump with low velocity dispersion $\\simeq 0.5$ km s$^{-1}$. K03 concluded that the secondary peak in UMi is a long-lived structure, surviving in phase-space because the underlying gravitational potential is close to harmonic. In the standard dark matter paradigm, this implies that the dark matter halo in UMi must have a large core because, if the dark halo has a central density cusp, the clump should have diluted in $\\sim 1$ Gyr. Even if dSph galaxies are influenced to some degree by the tidal forces exerted by the Milky Way, it is unlikely that such a large core may have a tidal origin (e.g., Stoehr et al.~2002; Hayashi et al.~2003; Pe\\~narrubia et al.~2008). It is worthwhile exploring if the competing MOND scenario can explain the survival of cold substructures in dSph galaxies in a natural way. The question that arises is whether the MOND gravitational potential can mimic the potential of a dark halo with a core in order to explain the very longevity of the dynamical fossil in UMi. In other words, can the inference of cored haloes in dSphs be accomodated naturally within MOND?  This paper is organized as follows. In section \\ref{sec:tidalradius}, we give a general statement of the problem and outline some key analytic results for tidal dissolution of satellite systems. In section \\ref{sec:UMidm}, we briefly describe the observational properties of UMi dwarf and its clump. Some important issues regarding the assumptions and approximations made to study the evolution of the clump in MOND framework are given in section \\ref{sec:approximations}. Section \\ref{sec:survival} presents results on the evolution of a stellar clump in MOND. In section \\ref{sec:alternatives},  we discuss alternative scenarios to account for the persistence of kinematically cold substructure in dSph galaxies. Finally,  our conclusions are summarized in section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  Observations of the rotation curves of spiral galaxies strongly support a one-to-one relation between gravity at any radius and the enclosed baryonic mass. Due to this empirical relation, modified gravity theories like MOND are able to account rather successfully for the amplitude and the shape of the rotation curves.  There are some indirect phenomena that suggest that this empirical relation may break down at scales of dSph galaxies (e.g., Gilmore et al.~2007, and references therein), challenging the interpretation of a modification of gravity as a substitute for dark matter.   Localized regions with enhanced stellar density and, where data permit, extremely cold kinematics have been detected in some dSph galaxies (e.g., Olszewski \\& Aaronson 1985; K03; Coleman et al.~2004; Walker et al.~2006). In the framework of CDM, K03 show that adopting a cored halo profile can preserve the UMi's clump incorrupted for a Hubble time. This cored halo can fit the observed stellar velocity dispersion and the persistence of the clump. However, although we are still unable to make robust predictions about how the dark matter distribution changes in the process of galaxy formation when the physics of baryons are included, the formation of a large core in a dark matter dominated galaxy, provides a hard challenge for $\\Lambda$CDM. Mashchenko et al.~(2008) showed that energy feedback in dwarf galaxies drives bulk gas motions and gravitational potential fluctuations large enough to turn the cusp into a flat core (see also Governato et al.~2009). They find that for a Fornax sized galaxy, the dark matter core has an average density of $0.2 M_{\\odot}$ pc$^{-3}$ at redshift $z=5.2$, and the density decays a factor $\\sqrt{2}$ in $\\sim 300$ pc. Still, such a core would not be enough to avoid orbital spreading of the clump.  In the particular case of UMi dSph galaxy, the density excess around the secondary peak cannot be the remnant of a merger with UMi of a smaller, gas-rich system because the stars have the same properties in terms of color and magnitude as the body of the UMi population (Kleyna et al.~1998). One might wonder whether the density peak could instead be a projection effect and that what we are seeing is a cold, low-density tidal tail. Numerical experiments by Read et al.~(2006) have shown that this scenario is very unlikely.  Another possibility is that the clump is a portion of the stationary debris of a disrupted globular cluster. This model is not implausible but cannot be proven yet.  A more drastic alternative scenario consists in interpreting the survival of substructures as an internal inconsistency of $\\Lambda$CDM. We have explored if MOND can explain the longevity of unbound clumps in dSph galaxies.  Following K03, we have assumed that this overdensity structure is a disrupted stellar cluster and simulate its evolution in the gravitational potential derived in MOND. Whichever the form of the orbit, the clump is tidally disrupted within $1.5$ Gyr, even if clump's self-gravity is included. One can decrease UMi mass by adopting a smaller $\\Upsilon_{\\star}$, slowing the process of orbital mixing because the crossing time for the clump becomes larger. However,  even assuming a mass-to-light ratio of $0.8 M_{\\odot}/L_{\\odot}^{V}$, the clump is disrupted in $<2.5$ Gyr. Our conclusion that tidal forces should have disrupted the clump appears robust for the adopted value of the stellar mass-to-light ratio of UMi.  The external field acceleration felt by UMi from the Milky Way depends on the galactocentric distance of UMi, which is time dependent if the orbit of UMi around the Galactic halo is elongated. However, the temporal variation of the EFE can hardly boost the longevity of the clump.    In the absence of any alternative model to explain the origin of the cold clump, we conclude that it is challenging for both $\\Lambda$CDM and MOND to explain the nature and dynamics of the clump in UMi. Another alternative of gravity suggested by Moffat (2005) is the so-called Modified Gravity (MOG). It is a fully covariant theory which predicts a Yukawa-like modification of Newton's law. For a system with a baryonic mass of a few times $10^{6} M_{\\odot}$ like UMi dSph galaxy, MOG predicts little or no observable deviation from Newtonian gravity at galactic distances of $\\sim 200$ pc (e.g., Moffat \\& Toth 2008). Therefore, explaining the internal dynamics of dSph galaxies is problematic without advocating dark matter, weakening the appeal of the MOG model. Other authors achieve to find modified theories of gravity which seem to reproduce the rotation curves of galaxies (e.g., Capozziello et al.~2007). Nevertheless, an analysis of the dynamics of dSph galaxies in these theories is still missing."
    },
    "1005/1005.5069_arXiv.txt": {
        "abstract": "The CMB maps obtained by observations always possess domains which have to be masked due to severe uncertainties with respect to the genuine CMB signal. Cosmological analyses ideally use full CMB maps in order to get e.\\,g.\\ the angular power spectrum. There are attempts to reconstruct the masked regions at least at low resolutions, i.\\,e.\\ at large angular scales, before a further analysis follows. In this paper, the quality of the reconstruction is investigated for the ILC (7yr) map as well as for 1000 CMB simulations of the $\\Lambda$CDM concordance model. The latter allows an error estimation for the reconstruction algorithm which reveals some drawbacks. The analysis points to errors of the order of a significant fraction of the mean temperature fluctuation of the CMB. The temperature 2-point correlation function $C(\\vartheta)$ is evaluated for different reconstructed sky maps which leads to the conclusion that it is safest to compute it on the cut-sky. ",
        "introduction": "The cosmic microwave background (CMB) provides one of the cornerstones of the cosmological concordance model. The statistical properties of  our cosmological models have to match those of the CMB in order to give an admissible model. Thus, it is of utmost importance to reliably extract the statistical properties of the CMB. A main obstacle is the foreground emission of our galaxy and of other sources which restrict the area of the sky available for a sufficiently clean CMB signal, i.\\,e.\\ the full sky CMB signal has to be masked. One path of statistical analysis leads to the Fourier space in which the CMB is decomposed with respect to spherical harmonics $Y_{lm}(\\hat n)$ where the masked sky leads to a coupling between the Fourier modes since no full sky CMB is available. In this paper, we do not delve into these difficulties, where upon an extensive literature exists, but instead follow the alternative path which allows an analysis directly in the pixel space, which in turn is more adapted to a masked sky. This analysis is based on the temperature two-point correlation function $C(\\vartheta)$, which is defined as \\begin{equation} \\label{Eq:C_theta} C(\\vartheta) \\; := \\; \\left< \\delta T(\\hat n) \\delta T(\\hat n')\\right> \\hspace{10pt} \\hbox{with} \\hspace{10pt} \\hat n \\cdot \\hat n' = \\cos\\vartheta \\hspace{10pt} , \\end{equation} where $\\delta T(\\hat n)$ is the temperature fluctuation in the direction of the unit vector $\\hat n$. The most direct way to deal with a mask is just to use only those pixels which are outside the mask. In this way it was discovered by the COBE team \\citep{Hinshaw_et_al_1996} that the correlation function $C(\\vartheta)$ possesses surprisingly low power at large angles $\\vartheta \\gtrsim 60^\\circ$.   A surprising observation is made by using the ILC map, which represents a full sky CMB map obtained by the WMAP team \\citep{Gold_et_al_2010}. Computing the correlation function $C(\\vartheta)$ using a mask leads to a correlation function having very low power at large scales, whereas using the complete ILC map leads to a correlation function which possesses higher large scale power being compatible with the concordance model \\citep{Spergel_et_al_2003}. One has to decide which correlation function $C(\\vartheta)$ corresponds to the true CMB sky: the one which is based on the safe pixels, i.\\,e.\\ those outside the mask, or the other one, which would imply that most of the large scale power is generated by those areas hidden by the galaxy, i.\\,e.\\ by those pixels which have experienced much larger corrections. This question has recently stimulated much discussions, e.\\,g.\\ \\citep{Copi_Huterer_Schwarz_Starkman_2006,Copi_Huterer_Schwarz_Starkman_2008,% Copi_Huterer_Schwarz_Starkman_2010,Hajian_2007,% Aurich_Janzer_Lustig_Steiner_2007,Aurich_Lustig_Steiner_2009,% Sarkar_Huterer_Copi_Starkman_Schwarz_2010,% Bennett_et_al_2010,% Efstathiou_Ma_Hanson_2009,Pontzen_Peiris_2010}.  \\cite{Efstathiou_Ma_Hanson_2009} emphasise that one has to start with the cut sky, but before the correlation function $C(\\vartheta)$ is computed the low-order multipoles have to be reconstructed. In this way, a stable result is obtained as long as the mask is not too large. The obtained correlation function is then the one with large power at large scales. For the details of the method see \\cite{deOliveira-Costa_Tegmark_2006,Bielewicz_Gorski_Banday_2004}. Here we summarise only the most important ingredients. The data vector $\\vec x$ containing only the pixel values outside the mask is related to the spherical harmonic coefficients $a_{lm}$ represented as $\\vec a$ by \\begin{equation} \\label{Eq:x_as_a} \\vec x \\; = \\; Y \\, \\vec a \\, + \\, \\vec n \\hspace{10pt} , \\end{equation} where $Y_{ij}$ denotes the corresponding values of $Y_{l_jm_j}(\\hat n_i)$ and $\\vec n$ the noise. In the low-order multipole reconstruction, only the multipoles with $l \\leq l_{\\hbox{\\scriptsize max}} = 10\\dots20$ are taken into account. The methods of reconstruction differ by the choice of a square matrix $A$ which determines the reconstructed $\\vec a\\,^r$ by \\begin{equation} \\label{Eq:ar_by_A} \\vec a\\,^r \\; = \\; (Y^T A Y)^{-1} \\, Y^T A \\vec x \\hspace{10pt} . \\end{equation} Setting the matrix $A$ equal to the unit matrix leads to the method of ``direct inversion''. To take the correlations between the pixels into account, one can choose the covariance matrix $A^{-1}= <\\vec x \\cdot \\vec x\\,^T>$ for a reconstruction up to $l_{\\hbox{\\scriptsize max}}$, which leads to the method used by \\cite{deOliveira-Costa_Tegmark_2006} with \\begin{equation} \\label{Eq:covariance_matrix} A_{ij}^{-1} \\; = \\; \\sum_{l=l_{\\hbox{\\scriptsize max}}+1}^{l_{\\hbox{\\scriptsize cut}}} \\frac{2l+1}{4\\pi} \\, P_l(\\hat n_i \\cdot \\hat n_j) \\, C_l \\end{equation} ignoring the noise contribution. The sum in (\\ref{Eq:covariance_matrix}) runs only over those multipole moments $C_l$, $l>l_{\\hbox{\\scriptsize max}}$, that are not to be reconstructed.  The reconstructed sky map up to $l_{\\hbox{\\scriptsize max}}$ is then obtained from the coefficients $\\vec a\\,^r$. \\cite{Efstathiou_Ma_Hanson_2009} compute the correlation function $C(\\vartheta)$ from such reconstructed maps for $l_{\\hbox{\\scriptsize max}}=5$ to $l_{\\hbox{\\scriptsize max}}=20$ by using the KQ85 or the KQ75 mask provided by the WMAP team, and it is found that these $C(\\vartheta)$ only show a negligible variation. This is interpreted as a sign for a stable method to compute the temperature correlation in the presence of masks. It should be emphasised that the WMAP team bases its 7 year investigation with respect to CMB anomalies \\citep{Bennett_et_al_2010} on the method of \\citep{Efstathiou_Ma_Hanson_2009}. This leads to the conclusion that the correlation function agrees well with the $\\Lambda$CDM concordance model and displays thus no anomalous behaviour.   The system of equations (\\ref{Eq:ar_by_A}) is over-determined, since there are much more pixel values than multipole moments with $l \\leq l_{\\hbox{\\scriptsize max}}$ as long as the mask and $l_{\\hbox{\\scriptsize max}}$ are not too large. Thus, there is the hope that the pixels outside the mask already determine the low-order multipoles. The caveat in the demonstration of \\citep{Efstathiou_Ma_Hanson_2009} is, however, that a large Gaussian smoothing of $10^\\circ$ is applied to the maps {\\it before} the cut is applied. Due to this smoothing there is a transfer of information from pixel values within the mask to those outside the mask. Since the system of equations (\\ref{Eq:ar_by_A}) is already over-determined, this additional information is readily extracted and reveals the temperature structure within the mask after the ``reconstruction''. It should be noted that even a downgrade in the HEALPix resolution can lead to an information transfer by carrying out the downgrade before the mask is applied.  \\begin{figure} \\begin{center} { \\begin{minipage}{11cm} \\hspace*{-20pt}\\includegraphics[width=9.0cm]{psplots/wmap_temperature_analysis_mask_r9_7yr_v4_small.eps} \\end{minipage} \\put(-315,67){(a) KQ85 (7yr) mask, $N_{\\hbox{\\scriptsize side}} = 512$} } \\vspace*{5pt} { \\begin{minipage}{11cm} \\hspace*{-20pt}\\includegraphics[width=9.0cm]{psplots/wmap_temperature_analysis_mask_r9_7yr_v4_s0.5_nside_16_small.eps} \\end{minipage} \\put(-315,67){(b) KQ85 (7yr) mask, $N_{\\hbox{\\scriptsize side}} = 16$, $x_{\\hbox{\\scriptsize th}}=0.5$ } } \\end{center} \\caption{\\label{Fig:KQ85_masks_nside_512_and_nside_16} The KQ85 mask of the WMAP 7 year data with a pixel resolution of $N_{\\hbox{\\scriptsize side}} = 512$ is displayed in panel \\ref{Fig:KQ85_masks_nside_512_and_nside_16}a. The masked region is pictured in black. Figure \\ref{Fig:KQ85_masks_nside_512_and_nside_16}b shows the KQ85 mask downgraded to a pixel resolution of $N_{\\hbox{\\scriptsize side}} =16$. All pixels with $x(i)\\le 0.5$, equation (\\ref{Eq:not_masked_pixel_downgr}), are considered as masked. } \\end{figure}  The two masks used in this paper are the KQ85 (7yr), shown in figure \\ref{Fig:KQ85_masks_nside_512_and_nside_16}a, and the KQ75 (7yr) mask which are available at the LAMBDA website. These are stored in the HEALPix \\citep{Gorski_Hivon_Banday_Wandelt_Hansen_Reinecke_Bartelmann_2005} format with a HEALPix resolution of $N_{\\hbox{\\scriptsize side}} = 512$. The reconstruction algorithm requires masks in lower resolutions of $N_{\\hbox{\\scriptsize side}} = 16$ or 32, and a downgrade has to be carried out. A single downgraded pixel $i$ contains $N_{\\hbox{\\scriptsize total}}$ pixels of the higher resolution map and from these only $N_{\\hbox{\\scriptsize nm}}$ pixel values are used in the averaging process, i.\\,e.\\ those that are {\\it not} masked in the higher resolution map, in order to compute the value of the downgraded pixel. If the ratio \\begin{equation} \\label{Eq:not_masked_pixel_downgr} x(i) \\; = \\; \\frac{N_{\\hbox{\\scriptsize nm}}} {N_{\\hbox{\\scriptsize total}}} \\end{equation} is larger than a given mask threshold $x_{\\hbox{\\scriptsize th}}\\in(0.0, 1.0)$, i.\\,e.\\ $x(i)>x_{\\hbox{\\scriptsize th}}$, the resulting pixel is {\\it not} masked, and otherwise it is masked. The result of the KQ85 (7yr) mask for $N_{\\hbox{\\scriptsize side}} =16$ and $x_{\\hbox{\\scriptsize th}}=0.5$ is shown in figure \\ref{Fig:KQ85_masks_nside_512_and_nside_16}b. It is obvious that the size of the masked domains in masks with $N_{\\hbox{\\scriptsize side}} <512$ depends on this mask threshold $x_{\\hbox{\\scriptsize th}}$.  \\begin{figure} \\begin{center} { \\begin{minipage}{11cm} \\hspace*{-20pt}\\includegraphics[width=9.0cm]{psplots/lcdm_7yr_wmap_bao_h0_nside_512_lmax_1024_fwhm_60.0arcmin_plus_fwhm_600.0arcmin_seed_126_sw_small.eps} \\end{minipage} \\put(-325,75){(a) $\\Lambda$CDM simulation A} } \\vspace*{10pt} { \\begin{minipage}{11cm} \\hspace*{-20pt}\\includegraphics[width=9.0cm]{psplots/lcdm_7yr_wmap_bao_h0_nside_512_lmax_1024_fwhm_60.0arcmin_plus_fwhm_600.0arcmin_seed_166_sw_small.eps} \\end{minipage} \\put(-325,75){(b) $\\Lambda$CDM  simulation B} } \\vspace*{10pt} { \\begin{minipage}{11cm} \\hspace*{-20pt}\\includegraphics[width=9.0cm]{psplots/lcdm_7yr_wmap_bao_h0_nside_512_lmax_1024_fwhm_60.0arcmin_seed_126_und_seed_166_maske_KQ85_7yr_600.0arcmin_sw_small.eps} \\end{minipage} \\put(-325,75){(c) $\\Lambda$CDM simulation A outside and B inside KQ85 mask} } \\end{center} \\caption{\\label{Fig:CutSky_Info_Input_maps} Figure \\ref{Fig:CutSky_Info_Input_maps}a shows the CMB simulation A of the $\\Lambda$CDM concordance model with a pixel resolution of $N_{\\hbox{\\scriptsize side}} = 512$ and an additional smoothing of $10^\\circ$. In figure \\ref{Fig:CutSky_Info_Input_maps}b a second map (CMB simulation B) using the same cosmological parameters is displayed. The map in figure \\ref{Fig:CutSky_Info_Input_maps}c results from the models used for figure \\ref{Fig:CutSky_Info_Input_maps}a and \\ref{Fig:CutSky_Info_Input_maps}b where the pixels outside the KQ85 mask are taken from simulation A and inside from simulation B at the resolution $N_{\\hbox{\\scriptsize side}} = 512$. Thereafter a Gaussian smoothing of $10^\\circ$ is applied. } \\end{figure}   \\begin{figure} \\begin{center} { \\begin{minipage}{10cm} \\hspace*{-20pt}\\includegraphics[width=9.0cm]{psplots/lcdm_7yr_wmap_bao_h0_nside_512_lmax_1024_fwhm_60.0arcmin_plus_fwhm_600.0arcmin_seed_126_kov_s0.5_lmax_10_sw_small.eps} \\end{minipage} \\put(-295,75){(a) $\\Lambda$CDM simulation A} } \\vspace*{10pt} { \\begin{minipage}{10cm} \\hspace*{-20pt}\\includegraphics[width=9.0cm]{psplots/lcdm_7yr_wmap_bao_h0_nside_512_lmax_1024_fwhm_60.0arcmin_plus_fwhm_600.0arcmin_seed_166_kov_s0.5_lmax_10_sw_small.eps} \\end{minipage} \\put(-295,75){(b) $\\Lambda$CDM  simulation B} } \\vspace*{10pt} { \\begin{minipage}{10cm} \\hspace*{-20pt}\\includegraphics[width=9.0cm]{psplots/lcdm_7yr_wmap_bao_h0_nside_512_lmax_1024_fwhm_60.0arcmin_seed_126_und_seed_166_maske_KQ85_7yr_600.0arcmin_kov_s0.5_lmax_10_sw_small.eps} \\end{minipage} \\put(-295,75){(c) $\\Lambda$CDM simulation A outside and B inside KQ85 mask} } \\end{center} \\caption{\\label{Fig:CutSky_Info_Extraction_lmax_10} The three reconstructed maps are displayed. The reconstructions are carried out for $l_{\\hbox{\\scriptsize max}}=10$. These maps result from the maps in figure \\ref{Fig:CutSky_Info_Input_maps}a, b and c by downgrading these maps to a pixel resolution of $N_{\\hbox{\\scriptsize side}} =16$ and thereafter applying the reconstruction algorithm (\\ref{Eq:ar_by_A}) to the data outside the mask shown in figure \\ref{Fig:KQ85_masks_nside_512_and_nside_16}b. Note that panel (a) and (c) should show the same sky map if the reconstruction algorithm would not use information from the masked region. However, there are significant difference especially near the Galactic plane. } \\end{figure}   In the next sections we will compare results of CMB simulations of the $\\Lambda$CDM concordance model with those of the ILC map. For this reason all maps of the $\\Lambda$CDM concordance model are produced in a FWHM resolution of $1^\\circ$ and $N_{\\hbox{\\scriptsize side}} =512$. But we also investigate these maps and the ILC map after an additional smoothing of e.\\,g. $10^\\circ$. In the following only the additional smoothing width is specified.   Let us now return to the information transfer caused by the smoothing procedure and/or by carrying out the downgrade. The following simple numerical experiment presented in figures \\ref{Fig:CutSky_Info_Input_maps} and \\ref{Fig:CutSky_Info_Extraction_lmax_10} reveals the information transfer. Figure \\ref{Fig:CutSky_Info_Input_maps}a displays the CMB simulation A of the $\\Lambda$CDM concordance model after a smoothing of $10^\\circ$. This map is downgraded from a pixel resolution of $N_{\\hbox{\\scriptsize side}} =512$ to $N_{\\hbox{\\scriptsize side}} =16$. Then the KQ85 mask in the pixel resolution $N_{\\hbox{\\scriptsize side}} =16$, shown in figure \\ref{Fig:KQ85_masks_nside_512_and_nside_16}b, is applied. For $l_{\\hbox{\\scriptsize max}}=10$ the reconstruction algorithm (\\ref{Eq:ar_by_A}) is used to obtain the reconstructed map shown in figure \\ref{Fig:CutSky_Info_Extraction_lmax_10}a. Both the original and the reconstructed map agree within the mask reasonably well. The same procedure is repeated for a second simulation B where figure \\ref{Fig:CutSky_Info_Input_maps}b shows the simulation and figure \\ref{Fig:CutSky_Info_Extraction_lmax_10}b the reconstruction. In the next step the pixels of simulation A are replaced at a pixel resolution of $N_{\\hbox{\\scriptsize side}} =512$ within the KQ85 mask, shown in figure \\ref{Fig:KQ85_masks_nside_512_and_nside_16}a, by those of simulation B shown in figure \\ref{Fig:CutSky_Info_Input_maps}b. After this replacement the smoothing of $10^\\circ$ is applied. The last step transfers now the ``wrong'' information to the pixels outside the mask. This smoothed map is shown in figure \\ref{Fig:CutSky_Info_Input_maps}c. Downgrading this map to $N_{\\hbox{\\scriptsize side}} =16$ provides the data outside the mask which are used for the reconstruction. If the reconstruction would not use the information within the mask, the reconstructed map of figure \\ref{Fig:CutSky_Info_Extraction_lmax_10}a should reappear. However, as revealed in figure \\ref{Fig:CutSky_Info_Extraction_lmax_10}c, the reconstruction algorithm generates within the mask the main structures of simulation B, which is displayed in figure \\ref{Fig:CutSky_Info_Input_maps}b. This clearly demonstrates the information transfer, so that one has to be careful in testing the reconstruction algorithm. This leads to the question whether the reconstruction can be carried out using only unsmoothed maps where no information about pixels within the mask is encoded outside. However, then stability difficulties arise as shown below.  \\begin{figure} \\begin{center} \\begin{minipage}{11cm} \\hspace*{-20pt}\\includegraphics[width=9.0cm]{psplots/wmap_ilc_7yr_v4_leak_600.0arcmin_small.eps} \\end{minipage} \\end{center} \\caption{\\label{Fig:wmap_ilc_7yr_v4_leak_600.0arcmin} The leak of the ILC temperatures from regions within the mask to those outside due to the smoothing process. } \\end{figure}  Another way to demonstrate the flow of information from regions within the mask to those outside the mask, is the following. At first the pixels of the ILC map outside the KQ85 mask are set to zero. Then a subsequent $10^\\circ$ smoothing shows how much information about the ILC pixel values inside the mask leaks to those regions outside. This leak of information is shown in figure \\ref{Fig:wmap_ilc_7yr_v4_leak_600.0arcmin}. It is obvious that the main structures of the ILC map within the mask appear close to the boundary outside the mask.  The paper is organised as follows. In section 2 the reconstruction errors of the CMB temperatures are evaluated outside and inside the masks using different smoothings and different resolutions. The reconstruction method using the covariance matrix is compared with the direct inversion method. In section 3 the influence of the reconstruction method onto the 2-point correlation function $C(\\vartheta)$ of the CMB is investigated where the focus is on large scales. An integrated measure of $C(\\vartheta)$ serves in section 4 as a further tool to demonstrate the drawbacks of the reconstruction methods. Finally, in section 5 we summarise our results. ",
        "conclusions": "In this paper we have investigated the reconstruction of pixel values within masks using inversion methods. We have compared the algorithm using the covariance matrix with the direct inversion method. The comparison shows that for the reconstruction of multipoles $l_{\\hbox{\\scriptsize max}}\\lesssim 8$ the algorithm using the covariance matrix gives slightly smaller errors than the direct inversion method. For this reason we focus on the reconstruction method which uses the covariance matrix.  An important point of our analysis is that the additional smoothing of $10^{\\circ}$ of the ILC map in \\cite{Efstathiou_Ma_Hanson_2009} transfers information from pixels inside the mask to pixels outside. This should not happen since the pixel values inside the mask are considered as insecure and should be ignored. The smoothing, however, leads to a quantitatively better reconstruction of the ILC map, if the masked domain is considered as containing the true pixel values which have to be recovered by the reconstruction. But the contribution of information from masked pixels makes a reconstructed map obtained from data with additional smoothing unacceptable in a realistic application.  The errors of the pixel values of the reconstructed map are compared with the mean temperature fluctuation $\\sigma_{\\hbox{\\scriptsize true}}(l_{\\hbox{\\scriptsize max}})$ defined in equation (\\ref{Eq:Ortsraum_Normierung}). In the case of the KQ75 (7yr) mask the reconstructed pixels inside the mask have errors larger than $\\sigma_{\\hbox{\\scriptsize true}}(l_{\\hbox{\\scriptsize max}})$ already for $l_{\\hbox{\\scriptsize max}} \\gtrsim 6$ such that the reconstruction is unusable as revealed by figure \\ref{Fig:Q_ilc_KQ75}. Even for smaller values of $l_{\\hbox{\\scriptsize max}}$ the errors are at least of the order of $\\frac 12\\sigma_{\\hbox{\\scriptsize true}}(l_{\\hbox{\\scriptsize max}})$. The culprit for this negative result is the size of the KQ75 (7yr) mask being too large for a reconstruction.   In the case of the smaller KQ85 (7yr) mask the reconstruction without additional smoothing leads for $l_{\\hbox{\\scriptsize max}} \\gtrsim 13$ to errors larger than $\\sigma_{\\hbox{\\scriptsize true}}(l_{\\hbox{\\scriptsize max}})$ such that a rough reconstruction can be done for $l_{\\hbox{\\scriptsize max}}$ around 10. However, even with this restriction, errors larger than $\\frac 12\\sigma_{\\hbox{\\scriptsize true}}(l_{\\hbox{\\scriptsize max}})$ occur, see figure \\ref{Fig:Q_ilc_KQ85}. This might be too large for cosmological applications. It is also checked that refining the resolution from $N_{\\hbox{\\scriptsize side}} = 16$, which is used in the above computations, to $N_{\\hbox{\\scriptsize side}} = 32$ does not improve the reconstruction.   Furthermore, the behaviour of the temperature 2-point correlation function $C(\\vartheta)$ is analysed with respect to the various reconstructed maps, i.\\,e.\\ with and without additional smoothing, for the ILC (7yr) map and for CMB simulations of the $\\Lambda$CDM concordance model. The correlation function $C(\\vartheta)$ is computed from the full ILC map, from the data outside the KQ85 mask, and from the reconstructed ILC maps which depend on $l_{\\hbox{\\scriptsize max}}$. In order to get a good approximation of the correlation function $C(\\vartheta)$ on large angular scales, the multipoles up to $l_{\\hbox{\\scriptsize max}} \\simeq 10$ are needed. The application of the large KQ75 mask is therefore excluded. The errors of the reconstructed 2-point correlation function are estimated by using 1000 CMB simulations of the $\\Lambda$CDM model. Since this error estimation is based on simulations which contain only the pure CMB signal, no further errors are taken into account e.\\,g.\\ resulting from residual foreground and detector noise. Therefore, the error estimation of the correlation function of the ILC map gives only a lower bound of the errors expected in a genuine application.   As discussed above the reconstructed ILC map with additional smoothing of 600 arcmin uses information from inside the mask. This information transfer leads in turn to underestimated errors in the 2-point correlation function. Furthermore, the similarity of $C(\\vartheta)$ computed from the full ILC map and from the reconstructed map is due to this information transfer and not the achievement of the reconstruction algorithm.   Without additional smoothing the ILC correlation function $C(\\vartheta)$, computed from the reconstructed map, does not match the one of the full map better than the one of the masked map within the estimated errors. This is caused by the large errors which increase with $l_{\\hbox{\\scriptsize max}}$. For $10 \\lesssim l_{\\hbox{\\scriptsize max}} \\lesssim 12$ the reconstructed 2-point correlation function is uncertain by 100$\\mu\\hbox{K}^2$ and for even larger values of $l_{\\hbox{\\scriptsize max}}$ useless as revealed by figure \\ref{Fig:correlation_function_ilc_KQ85_FWHM_000arcmin_nside_16}.   This uncertainty is also reflected in the $S(60^{\\circ})$ statistic which integrates the power of the 2-point correlation function on angular scales larger than $60^{\\circ}$. The $S(60^{\\circ})$ statistic of the reconstructed ILC map does not converge to a stable value for $l_{\\hbox{\\scriptsize max}} \\gtrsim 6$ as it is the case for the full and for the masked ILC map as displayed in figure \\ref{Fig:s_statistik_1000lcdm_und_ilc_KQ85_FWHM_000arcmin_nside_16}a. In addition, the mean value of the $S(60^{\\circ})$ statistic, calculated from 1000 reconstructed CMB sky maps, reveals the instability of the algorithm for $l_{\\hbox{\\scriptsize max}} \\gtrsim 12$, where the power strongly increases, see the dashed curve in figure \\ref{Fig:s_statistik_1000lcdm_und_ilc_KQ85_FWHM_000arcmin_nside_16}b. This demonstrates the inability to reconstruct the temperature correlations.   The conclusion of this paper is that by using a realistic mask for the WMAP data the reconstruction algorithm (\\ref{Eq:ar_by_A}) does not work well enough to obtain a prediction of the 2-point correlation function. For this reason a cosmological analysis of the WMAP data should use only data outside a mask, e\\,g.\\ the KQ85 (7yr) or KQ75 (7yr) mask. The Planck satellite measures the sky at more different frequencies than the WMAP satellite. This should allow a more secure reduction of the foreground in the CMB map. It is expected that the masked region in the Planck data is significantly smaller. It might be small enough in order to allow an acceptable reconstruction of the CMB within the mask which could provide a 2-point correlation function useable at large scales."
    },
    "1005/1005.2225_arXiv.txt": {
        "abstract": "\\ammonia\\ and \\m\\ are key molecules in astrochemical networks leading to the formation of more complex N- and O-bearing molecules, such as \\mc\\ and \\dime.  Despite a number of recent studies, little is known about their abundances in the solid state. This is particularly the case for low-mass protostars, for which only the launch of the \\spitzer\\ Space Telescope has permitted high sensitivity observations of the ices around these objects.  In this work, we investigate the $\\sim 8-10$~\\micron\\ region in the \\spitzer\\ IRS (InfraRed Spectrograph) spectra of 41 low-mass young stellar objects (YSOs). These data are part of a survey of interstellar ices in a sample of low-mass YSOs studied in earlier papers in this series.  We used both an empirical and a local continuum method to correct for the contribution from the 10~\\micron\\ silicate absorption in the recorded spectra.  In addition, we conducted a systematic laboratory study of \\ammonia- and \\m-containing ices to help interpret the astronomical spectra. We clearly detect a feature at $\\sim$9~\\micron\\ in 24 low-mass YSOs. Within the uncertainty in continuum determination, we identify this feature with the \\ammonia\\ $\\nu_2$ umbrella mode, and derive abundances with respect to water between $\\sim$2 and 15\\%.  Simultaneously, we also revisited the case of \\m\\ ice by studying the $\\nu_4$ C--O stretch mode of this molecule at $\\sim$9.7~\\micron\\ in 16 objects, yielding abundances consistent with those derived by \\citet{boogert-etal08} (hereafter paper I) based on a simultaneous 9.75 and 3.53~$\\mu$m data analysis. Our study indicates that \\ammonia\\ is present primarily in \\water-rich ices, but that in some cases, such ices are insufficient to explain the observed narrow FWHM. The laboratory data point to \\m\\ being in an almost pure methanol ice, or mixed mainly with CO or CO$_2$, consistent with its formation through hydrogenation on grains. Finally, we use our derived \\ammonia\\ abundances in combination with previously published abundances of other solid N-bearing species to find that up to 10--20~\\% of nitrogen is locked up in known ices. ",
        "introduction": "Ammonia and methanol are among the most ubiquitous and abundant (after H$_2$ and CO) molecules in space. Gaseous \\ammonia\\ and \\m\\ are found in a variety of environments such as infrared dark clouds, dense gas surrounding ultra-compact \\hii\\ regions, massive hot cores, hot corinos, and comets. Solid \\m\\ has been observed in the ices surrounding massive YSOs (e.g. \\citealt{schutte-etal91,dartois-etal99,gibb-etal04}) and more recently toward low-mass protostars \\citep{pontoppidan-etal03}. The presence of solid \\ammonia\\ has been claimed toward massive YSOs only \\citep{lacy-etal98,dartois-etal02,gibb-etal04,guertler02}, with the exception of a possible detection in the low-mass object IRAS 03445+3242 % \\citep{guertler02}. However, these detections are still controversial and ambiguous \\citep{taban-etal03}.  Both molecules are key participants in gas-grain chemical networks resulting in the formation of more complex N- and O-bearing molecules, such as \\mc\\ and \\dime\\ (e.g. \\citealt{rodgers+charnley01}). Moreover, UV processing of \\ammonia- and \\m-containing ices has been proposed as a way to produce amino-acids and other complex organic molecules (e.g., \\citealt{munozcaro+schutte03,bernstein-etal02,oberg-etal09}). In addition, the amount of \\ammonia\\ in ices has a direct impact on the content of ions such as NH$_4^+$ and OCN$^-$, which form reactive intermediates in solid-state chemical networks. A better knowledge of the \\ammonia\\ and \\m\\ content in interstellar ices will thus help to constrain chemical models and to gain a better understanding of the formation of more complex, prebiotic, molecules.  During the pre-stellar phase, \\ammonia\\ is known to freeze out on grains (if the core remains starless long enough -- \\citealt{lee-etal04}). Moreover, \\m\\ is known to have gas-phase abundances with respect to H$_2$ in hot cores/corinos that are much larger than in cold dense clouds: $\\sim (1-10)\\times 10^{-6}$ vs $\\leq 10^{-7}$, with the former values most likely representing evaporated ices in warm regions \\citep[e.g.][]{genzel-etal82,blake-etal87,federman-etal90}. Together, these findings suggest that ices are an important reservoir of \\ammonia\\ and \\m\\ and that prominent features should be seen in the absorption spectra toward high- and low-mass protostars. Unfortunately, as summarized in Table~\\ref{tab:features}, \\ammonia\\ and \\m\\ bands, with the exception of the 3.53~\\micron\\ \\m\\ feature, are often blended with deep water and/or silicate absorptions, complicating unambiguous identifications and column density measurements. This is particularly true for \\ammonia\\ whose abundance determination based on the presence of an ammonium hydrate feature at 3.47~\\micron\\ remains controversial (e.g. \\citealt{dartois+dhendecourt01}). Nonetheless, it is important to use all available constraints to accurately determine the abundances of these two molecules. Despite the overlap with the 10~\\micron\\ silicate (Si--O stretch) feature, the \\ammonia\\ $\\nu_2$ umbrella mode at $\\sim$9~\\micron\\ ($\\sim$1110~\\wn) offers a strong intrinsic absorption cross section and appears as the most promising feature to determine the abundance of this species in the solid phase. Moreover, the \\m\\ $\\nu_4$ C--O stretch at $\\sim$9.7~\\micron\\ ($\\sim$1030~\\wn) provides a good check on the validity of the different methods we will use to subtract the 10-\\micron\\ silicate absorption, since the abundance of this molecule has been accurately determined previously from both the 3.53 and 9.75~\\micron\\ features (see Paper I).   \\begin{deluxetable}{lccp{6cm}} \\tablewidth{0pt} \\tablecaption{Selected near- and mid-infrared features of \\ammonia\\ and \\m. \\label{tab:features}} \\tablehead{Mode\t& $\\lambda$ (\\micron)\t& $\\bar{\\nu}$ (cm$^{-1}$) & Problem} \\startdata \\sidehead{\\ammonia\\ features:} $\\nu_3$ N--H stretch\t& \\phn 2.96\t& 3375\t& Blended with \\water\\ (O--H stretch, 3.05~\\micron/3275~\\wn) \\\\ $\\nu_4$ H--N--H bend\t& \\phn 6.16\t& 1624\t& Blended with \\water\\ (H--O--H bend, 5.99~\\micron/1670~\\wn), HCOOH\\\\ {\\bf \\boldmath{$\\nu_2$} umbrella}\t\t& {\\bf \\phn 9.00}\t& {\\bf 1110}\t& {\\bf Blended with silicate} \\\\ \\hline \\sidehead{\\m\\ features:} $\\nu_2$ C--H stretch\t& \\phn 3.53\t& 2827\t& --\\\\ $\\nu_6$ \\& $\\nu_3$ --CH$_3$ deformation\t& \\phn 6.85\t& 1460\t& Blended (e.g. with NH$_4^+$)\\\\ $\\nu_7$ --CH$_3$ rock\t& \\phn 8.87\t& 1128\t& Weak; blended with silicate\\\\ {\\bf \\boldmath{$\\nu_4$} C--O stretch}\t& {\\bf \\phn9.75}\t& {\\bf 1026}\t& {\\bf Blended with silicate} \\\\ Torsion\t\t& 14.39\t\t& \\phn695\t& Blended with \\water\\ libration mode\\\\ \\enddata \\tablecomments{The bold-faced lines indicate the features studied here.} \\tablecomments{The nomenclature for the NH$_3$ and CH$_3$OH vibrational modes are adopted from \\citet{herzberg45}. } \\end{deluxetable}   More detailed spectroscopic information is particularly interesting for low-mass protostars as the ice composition reflects the conditions during the formation of Sun-like stars.  Such detections have only become possible with \\spitzer, whose sensitivity is necessary to observe low luminosity objects even in the nearest star-forming clouds. The spectral resolution of the \\spitzer\\ Infrared Spectrograph (IRS, \\citealt{houck-etal04}) of $\\Delta \\lambda/\\lambda \\sim 100$ in this wavelength range is comparable to that of the Infrared Space Observatory (ISO) PHOT-S instrument but lower than that of the ISO-SWS and other instruments used to identify solid \\ammonia\\ toward high luminosity sources. The spectral appearance of ice absorption features, such as band shape, band position and integrated band strength, is rather sensitive to the molecular environment. Thus, the interpretation of the astronomical spectra should be supported by a systematic laboratory study of interstellar ice analogues containing \\ammonia\\ and \\m. Changes in the lattice geometry and physical conditions of an ice are directly reflected by variations in these spectral properties. In the laboratory, it is possible to record dependencies over a wide range of astrophysically relevant parameters, most obviously ice composition, mixing ratios, and temperature. Such laboratory data exist for pure and some H$_2$O-rich NH$_3$- and CH$_3$OH-containing ices \\citep[e.g.,][]{dhendecourt+allamandola86,hudgins-etal93,kerkhof-etal99,taban-etal03}, but a systematic study and comparison with observational spectra is lacking.  In principle, the molecular environment also provides information on the formation pathway of the molecule. For example, NH$_3$ ice is expected to form simultaneously with H$_2$O and CH$_4$ ice in the early, low-density molecular cloud phase from hydrogenation of N atoms \\citep[e.g.,][]{tielens+hagen82}.  In contrast, solid CH$_3$OH is thought to result primarily from hydrogenation of solid CO, a process which has been confirmed to be rapid at low temperatures in several laboratory experiments \\citep[e.g.][]{watanabe+kouchi02,hidaka-etal04,fuchs-etal09}. A separate, water-poor layer of CO ice is often found on top of the water-rich ice layer in low-mass star-forming regions due to the `catastrophic' freeze-out of gas-phase CO at high densities \\citep{pontoppidan-etal03,pontoppidan06}. Hydrogenation of this CO layer should lead to a nearly pure CH$_3$OH ice layer \\citep[e.g.,][]{cuppen-etal09}, which will have a different spectroscopic signature from that of CH$_3$OH embedded in a water-rich matrix.  The latter signature would be expected if CH$_3$OH ice were formed by hydrogenation of CO in a water-rich environment or by photoprocessing of H$_2$O:CO ice mixtures, another proposed route \\citep[e.g.,][]{moore+hudson98}.    Here, we present \\spitzer\\ spectra between 5 and 35~\\micron\\ of ices surrounding 41 low-mass protostars, focusing on the $\\sim 8-10$~\\micron\\ region that contains the $\\nu_2$ umbrella and $\\nu_4$ C--O stretch modes of \\ammonia\\ and \\m, respectively. This work is the fourth paper in a series of ice studies \\citep{boogert-etal08,pontoppidan-etal08-co2,oberg-etal08-ch4} carried out in the context of the \\spitzer\\ Legacy Program ``From Molecular Cores to Planet-Forming Disks'' (``c2d''; \\citealt{evans-etal03}). In Section~\\ref{sec:astro}, we carry out the analysis of the \\spitzer\\ data in $8-10~\\micron$ range. In Section~\\ref{sec:lab}, we present the laboratory data specifically obtained to help interpret the data that are discussed in Section~\\ref{sec:comparison}. Finally, we conclude  in Section~\\ref{sec:ccl} with a short discussion of the joint astronomy-laboratory work (including the overall continuum determination). ",
        "conclusions": "\\label{sec:ccl}  We have analyzed in detail the 8-10~\\micron\\ range of the spectra of 41 low-mass YSOs obtained with \\spitzer\\ and presented in \\citet{boogert-etal08}.  The sources are categorized into three types: straight, curved and rising 8~\\micron\\ silicate wings, and for each category template sources with little or no absorption from ices around 9-10~\\micron\\ have been determined.  This has led to two ways of subtracting the contribution from the 10~\\micron\\ silicate absorption: first, by determining a local continuum, and second, by scaling the templates to the optical depth at 9.7~\\micron.  The two methods give consistent band positions of the NH$_3$ features, but the resulting widths can be up to a factor of two larger using the template continuum method. Taking into account the uncertainty in continuum determination, NH$_3$ ice is most likely detected in 24 of the 41 sources with abundances of $\\sim$2 to 15~\\% w.r.t. \\water, with an average abundance of 5.5$\\pm$2.0~\\%. These abundances have estimated uncertainties up to a factor of two and are not inconsistent with other features in the 3 and 6~\\micron\\ ranges.  CH$_3$OH is often detected as well, but the NH$_3$/CH$_3$OH abundance ratio changes strongly from source to source. Our inferred CH$_3$OH column densities are consistent with the values derived in paper I.  Targeted laboratory experiments have been carried out to characterize the NH$_3$ and CH$_3$OH profiles (position, FWHM, integrated absorbance).  Comparison with the observational data shows reasonable agreement (within $\\sim$1\\%) for the position of the NH$_3$ feature in H$_2$O-rich ices, but the observed widths are systematically smaller than the laboratory ones for nearly all sources. The silicate template continuum method gives widths that come closest to the laboratory values.  This difference in width (i.e. widths derived from astronomical spectra smaller than those in the laboratory spectra) suggests that the NH$_3$ abundances determined here may be on the low side.  The CH$_3$OH profile is most consistent with a significant fraction of the CH$_3$OH in a relatively pure or CO-rich phase, consistent with its formation by the hydrogenation of CO ice. In contrast, the most likely formation route of NH$_3$ ice remains hydrogenation of atomic N together with water ice formation in a relatively low density molecular phase. Finally, the nitrogen budget indicates that up to 10 to 20~\\% of nitrogen is locked up in known ices."
    },
    "1005/1005.5319_arXiv.txt": {
        "abstract": "{Recently, wavelets and $R / S$ analysis have been used as statistical tools to characterize the optical flickering of cataclysmic variables.} {Here we present the first comprehensive study of the statistical properties of X-ray flickering of cataclysmic variables in order to link them with physical parameters.} {We analyzed a sample of 97 X-ray light curves of 75 objects of all classes observed with the \\textit{XMM}-Newton space telescope. By using the wavelets analysis, each light curve has been characterized by two parameters, $\\alpha$ and $\\Sigma$, that describe the energy distribution of flickering on different timescales and the strength at a given timescale, respectively. We also used the $R / S$ analysis to determine the Hurst exponent of each light curve and define their degree of stochastic memory in time.} {The X-ray flickering is typically composed of long time scale events $(1.5 \\la \\alpha \\la 3)$, with very similar strengths in all the subtypes of cataclysmic variables $(-3 \\la \\Sigma \\la -1.5)$. The X-ray data are distributed in a much smaller area of the $\\alpha -\\Sigma$ parameter space with respect to those obtained with optical light curves. The tendency of the optical flickering in magnetic systems to show higher $\\Sigma$ values than the non-magnetic systems is not encountered in the X-rays. The Hurst exponents estimated for all light curves of the sample are larger than those found in the visible, with a peak at 0.82. In particular, we do not obtain values lower than 0.5. The X-ray flickering presents a persistent memory in time, which seems to be stronger in objects containing magnetic white dwarf primaries.} {The similarity of the X-ray flickering in objects of different classes together with the predominance of a persistent stochastic behavior can be explained it terms of magnetically-driven accretion processes acting in a considerable fraction of the analyzed objects.} ",
        "introduction": "Cataclysmic variables (CVs) are binary systems in which a late-type secondary star transfers matter onto a white dwarf (WD) primary. The accretion configuration is different depending on the magnetic field strength of the WD. In non-magnetic systems the accretion disk extends down to the WD surface, while a truncated disk may form in moderately magnetic CVs or may not be present at all in strongly magnetized systems (see \\citealt{warner95} for a review).  Owing to the gravitational potential of the compact primary, the accretion of matter produces a non-negligible flux of X-rays. The mechanism responsible for this emission in non-magnetic CVs during quiescence is the shock heating that acts in the boundary layer between the accretion disk and the WD surface \\citep{patray85b,patray85a}, while in magnetic systems X-rays are emitted by a standing shock above the magnetic poles of the WD \\citep{aizu73}. In the latter case the X-ray flux is higher, which makes the magnetic systems the brightest X-ray CVs.  Cataclysmic variables are subdivided into three main classes: novae (classical and recurrent), dwarf novae (DNe) and nova-like (NL) systems. In addition, there is a class of objects closely related to novae, the so-called super-soft X-ray sources (SSSs), whose members are characterized by a prominent soft spectral component due to thermonuclear burning at the WD surface \\citep[see e.g.][]{orio95}. A parallel classification can be made considering the strength of the magnetic field $B$ of the primary. In this case we have the non-magnetic systems, the intermediate polars (IPs, $B \\sim 5-10$~MG) and the polars ($B \\ga 10$~MG).  The X-ray light curves of CVs may exhibit a variety of modulations. Periodic coherent modulations due to occultations of the emitting regions can be produced by the orbital motion or by the rotation of the WD. For magnetic systems periodic modulations can also be ascribed to absorption effects produced in the magnetically-confined accretion flow [``columns'' in polars \\citep{cropper90}, ``curtains'' in IPs \\citep{rosetal88}]. Besides this persistent variability, CVs may also show quasi-periodic oscillations (QPOs) in both the soft and hard X-ray bands as well as completely stochastic brightness variations, which are usually identified with the term ``flickering'' \\citep[see][and references therein]{kuuetal06}.  Flickering is generally constituted by a sequence of random flares with typical timescales ranging from a few seconds to a few minutes. This phenomenon has been observed in the X-ray light curves of CVs of all classes. For example, observations with the \\textit{HEAO}-1 satellite allowed the discovery of soft X-ray flickering in the polar prototype \\object{AM~Her} \\citep{szketal80}, as well as hard X-ray aperiodic variability in the DN \\object{SS~Cyg} during quiescence \\citep{coretal84}. Flickering at timescales of $\\sim 10$~s was detected in the NL \\object{TT~Ari} using both \\textit{Einstein} \\citep{jenetal83} and \\textit{ASCA} \\citep{baykiz96} data. Rapid aperiodic fluctuations were discovered in the light curves of a bunch of CVs also using \\textit{ROSAT} observations \\citep{holetal94,roselta95,bucetal98}. A common observational feature is that X-ray flickering shows a continuous power law frequency spectrum, as it does in the visible region. A correlation between the time scale of the X-ray flickering activity and that detected in simultaneous optical observations was found in a number of CVs of different classes, e.g. the IP system \\object{V795~Her} \\citep{roselta95}, the polar \\object{EF~Eri} \\citep{watetal87}, the \\object{VY~Scl} star TT~Ari \\citep{jenetal83}, the old nova \\object{V603~Aql} \\citep{dreetal83} and in the DNe SS~Cyg and \\object{U~Gem} \\citep{coretal84}. This observational evidence has been commonly explained as the result of reprocessing of the X-ray flickering energy.  The properties of flickering have been studied for a long time, especially in the visible region of the electromagnetic spectrum. However, the origin of optical flickering is still uncertain, although there is plenty of evidence that it has to be related to the accretion process. It is very likely that the location of its source should be restricted either to regions very close to the WD, like the innermost part of the disk, or the hot spot (see \\citealt{bruch92} for a thorough discussion). The connection between X-ray flickering and the accretion onto the WD, instead, is evident. For highly magnetic CVs, the rapid soft X-ray flares often detected in their light curves could been explained as the result of a bombardment of the WD surface by random inhomogeneous structures (``blobs'') present in the accretion streams \\citep{kuipri82}. For instance, the features observed in the light curves of the polar \\object{V1309~Ori} \\citep{demetal98,schetal05} and of the asynchronous polar \\object{BY~Cam} \\citep{ramcro02}, although they are quite peculiar objects, could be explained with this hypothesis. Also the flares observed in the X-ray light curve of the old nova \\object{GK~Per}, which is an IP, have been interpreted as an indirect indication of blobby accretion \\citep{vrieetal05}.  In general, it is often possible to obtain some information about the flickering from the observations of a single object. However, as this is a stochastic process in time, it can be better studied by means of its statistical properties. \\citet[][hereafter FB98]{fribru98} carried out a statistical analysis of the optical flickering of a large sample of CVs based on the wavelet analysis \\citep{daubechies92} of their light curves. They represented the properties of flickering in a two-dimensional parameter space and showed that CVs of different classes tend to occupy different regions of this space. Later \\citet[][hereafter TDB09]{tametal09} used the $R / S$ rescaled range analysis \\citep{hurst51,huretal65} for the first time as a complementary tool to characterize the degree of persistence/anti-persistence of the white light flickering of the IP-class CV \\object{V709~Cas}. Motivated by the interesting results presented in these works, we have therefore used the same statistical techniques to study the X-ray flickering properties of CVs as a whole as well as a function of their individual classes to link them with physical parameters. ",
        "conclusions": "We have used the wavelets and the $R / S$ analysis as statistical tools to characterize the flickering in a sample of 97 X-ray light curves of CVs of all classes.  In general, the distribution of the X-ray data points in the $\\alpha -\\Sigma$ parameter space is almost independent of the class and, therefore, quite different from that in the visible. The $\\alpha$ parameters are all contained in the range $1.5 \\la \\alpha \\la 3$. This implies that the dissipation of the flickering energy occurs typically in long time scale flares. Moreover, the strength of the X-ray flickering is found to vary in the region $-3 \\la \\Sigma \\la -1.5$, which is much smaller than that found in the optical region. There is no evidence for magnetic NL systems to have higher $\\Sigma$ values than the non-magnetic systems. The different behavior in the visible region and in the X-rays could be explained assuming that part of the optical flickering in system harboring an accretion disk could be originated by other slower components such as trailing waves.  All objects in our sample show values of the Hurst exponent higher than 0.5, with a peak of the distribution at $H = 0.82$. The predominance of a persistent stochastic behavior can be ascribed to a magnetically-driven accretion process acting in the majority of the CVs considered here. In particular, the result obtained for magnetic CVs clearly depicts the behavior of a plasma trapped by the magnetic field in the accretion streams.  However, it must be pointed out that the sample of the X-ray light curves analyzed here is still not large enough to allow us to draw firm conclusions. Therefore, the analysis of further data collected in future X-ray observations will permit us to overcome possible selection effects and to better define the statistical properties of the flickering. On the other hand, a larger sample of polar systems analyzed in the optical region would allow us to verify our suggestion that the observed properties of X-ray flickering are intimately related to magnetically-driven accretion processes.  As a final remark, it would have been interesting to use the method proposed by TBD09 to calculate the Hurst exponent from the $\\alpha$ parameter (or vice-versa) and quantitatively demonstrate the impact of that procedure. However, it was not possible to do it with the X-ray data because for the vast majority the of the objects in the sample the $R / S$ analysis has been typically stopped at a timescale comparable to the second or third ``octave'' of the corresponding scalegrams. For this reason, the $H$ values obtained with the two methods are not comparable because either they refer to different timescales maybe including multiple Hurst exponents, or the $\\alpha$ parameter is obtained by linearly fitting just two or three points, which is not reliable."
    },
    "1005/1005.3824_arXiv.txt": {
        "abstract": "{The Sunyaev--Zel'dovich (SZ) effect is a spectral distortion of the cosmic microwave background as observed through the hot plasma in galaxy clusters.  This distortion is a decrement in the CMB intensity for $\\lambda > 1.3 \\,$mm, an increment at shorter wavelengths, and small again by $\\lambda \\sim 250 \\, \\mu$m.  As part of the \\textit{Herschel} Lensing Survey (HLS) we have mapped \\bulletc\\ (the Bullet cluster) with SPIRE with bands centered at $250$, $350$ and $500 \\, \\mu$m and have detected the SZ effect at the two longest wavelengths.  The measured SZ effect increment central intensities are $\\Delta I_{0} = 0.097 \\pm 0.019 \\,$MJy sr$^{-1}$ at $350 \\, \\mu$m and $\\Delta I_{0} = 0.268 \\pm 0.031 \\,$MJy sr$^{-1}$ at $500 \\, \\mu$m, consistent with the SZ effect spectrum derived from previous measurements at $2 \\,$mm. No other diffuse emission is detected.  The presence of the finite temperature SZ effect correction is preferred by the SPIRE data at a significance of $2.1 \\sigma$, opening the possibility that the relativistic SZ effect correction can be constrained by SPIRE in a sample of clusters.  The results presented here have important ramifications for both sub-mm measurements of galaxy clusters and blank field surveys with SPIRE.} ",
        "introduction": "The Sunyaev-Zel'dovich (SZ) effect is a distortion of the spectral shape of the cosmic microwave background (CMB) due to inverse Compton scattering in the ubiquitous, hot ($T_{\\mathrm{e}} \\sim 10^{7} \\,$K) intracluster medium (ICM) of galaxy clusters \\citep{Sunyaev1972}.  The canonical thermal SZ spectrum is a decrement in the brightness of the CMB as measured through galaxy clusters at mm wavelengths and an increment at sub-mm wavelengths which passes though a null at $\\lambda \\approx 1.3 \\,$mm.  To correctly describe the SZ spectral distortion when relativistic electrons are present or the cluster is moving with respect to the CMB additional correction terms, usually termed ``relativistic'' (or ``finite temperature'') corrections and ``kinetic'' SZ effect, are required.  Measurement of these corrections is only possible using observations at multiple wavelengths, and is expedited by measurement at wavelengths where the expected modifications to the thermal SZ effect are largest.  In the case of the finite temperature corrections, the largest changes expected in the SZ increment are at wavelengths shorter than $1 \\, \\mu$m.  Decrements in emission are rare astrophysically and can be ascribed to the SZ effect with little ambiguity; this has lead to measurements of the SZ effect at $\\lambda \\gtrsim 2 \\,$mm becoming almost routine.  In comparison, measurements of the SZ effect increment are complicated by the presence of the dusty, high redshift galaxies which constitute the sub-mm cosmic background.  Additionally, the individual sources comprising the sub-mm background are gravitationally lensed by galaxy clusters, the effect of which is to preferentially correlate increases in sub-mm emission with clustering.  This correlation makes unambiguous detection of the SZ increment difficult, though successful measurements do exist (\\citealt{Lamarre1998}, \\citealt{Komatsu1999}, \\citealt{Zemcov2003}, \\citealt{Zemcov2007}, \\citealt{Nord2009}). Moreover, the presence of the sub-mm background may contaminate measurements of the SZ effect for $\\lambda > 1 \\,$mm in less massive galaxy clusters \\citep{Aghanim2005}.  A better understanding of the sub-mm emission associated with galaxy clusters is required.  Heretofore, systematic far infrared (FIR) surveys of many galaxy clusters to large radii have been technically challenging so a complete census of sub-mm emission from clusters has been difficult to obtain.  The advent of SPIRE \\citep{Griffin2010} on \\textit{Herschel} \\citep{Pilbratt2010} has, for the first time, provided the capability to make deep maps of clusters to large angles on the sky and to use colour information to separate the different sources of sub-mm emission present in galaxy clusters.  In addition to gravitationally lensed background sources \\citep{Rex2010}, emission in clusters above the confused sub-mm background may also comprise emission from galaxies in the cluster itself \\citep{Rawle2010}, as well as truly diffuse emission from the SZ effect and possibly even cold dust in the ICM.  SPIRE's ability to separate sources of emission based both on spatial and spectral information allows the demographics of the sub-mm emission to be measured.  In this paper, we use deep SPIRE maps of the $z=0.3$ Bullet cluster (\\bulletc) taken as part of the \\textit{Herschel} Lensing Survey (HLS, P.I.~Egami) at $250$, $350$, and $500 \\, \\mu$m with $18$, $25$ and $36$ arcsec resolution to measure the SZ effect and constrain other diffuse emission associated with the cluster. ",
        "conclusions": "\\label{S:conclusion}  The possibility of detecting the SZ effect significantly shortward of its positive peak is a testament to the extraordinary capabilities of SPIRE.  However, works like \\citet{Zemcov2007} show that in large galaxy cluster survey samples, significant contamination to the SZ signal from bright, gravitationally lensed background sources is common.  Though the properties of this particular cluster have not precluded measurement of the SZ effect -- \\bulletc\\ has fewer bright, lensed sources close to its SZ effect center than typical clusters which have been observed in the sub-mm and is relatively bright and broad in the SZ effect -- based on this single example it is difficult to determine whether gravitational lensing of the sub-mm background will make measurements similar to this one more challenging in more typical clusters at these and other wavelengths.  Data from surveys like the HLS will allow us to understand whether typical clusters are suitable for SZ effect increment detection, and how the lensed sub-mm background will affect measurements at other wavelengths.  The possibility of diffuse emission from cold dust in the ICM has been discussed in the past (\\textit{e.g.} \\citealt{Stickel2002}).  It is expected that, due to sputtering by energetic photons, such dust would have a very short lifetime in the ICM environment \\citep{Draine1979}. Based on the radial averages of the $250 \\, \\mu$m source subtracted maps, where for ICM dust with reasonable temperatures the brightest thermal emission would occur, we find no evidence for this type of diffuse emission in this cluster.  Because in any reasonable scenario such dust emission would be faint in a $z > 0.1$ cluster, we expect that targeted searches of local clusters will be more successful for this science, though the HLS and similar surveys can provide useful constraints.  Given the statistical uncertainties, the $350 \\, \\mu$m SZ effect increment measured here is slightly less than $2 \\sigma$ larger than would be expected from the $2 \\,$mm measurements.  This may point to residual problems with the source subtraction procedure.  More data will allow tuning of the confused sub-mm background removal algorithm and checks on whether biases arising from poor background removal are endemic in a large cluster sample.  Given the presence of the SZ effect in clusters, it seems that care must be taken with photometry of $500 \\, \\mu$m sources within an arcminute or so of the cluster center; such sources will be positioned on a diffuse background so their fluxes will be biased by a small amount \\citep{Rex2010}.  As the peak of the FIR emission of dusty sources is redshifted to progressively longer wavelengths, high redshift sources are expected to have exceptionally red spectral energy distributions.  The results of this work show that care must be taken when searching for sources based on their presence in the $500 \\, \\mu$m band alone; galaxy clusters whose SZ effects are relatively bright and compact could well masquerade as such sources.  Due to the well known problem of determining counterparts to sub-mm sources at other wavelengths, such SZ effect contamination may not be immediately obvious, particularly as cluster fields are crowded and many possible counterparts may be present.  As a corollary, searching for very red sources in confusion limited SPIRE surveys may turn up compact clusters based on the presence of strong $500 \\, \\mu$m emission.  The search for such extremely red SPIRE sources in the HLS and other programmes is underway now."
    },
    "1005/1005.3223_arXiv.txt": {
        "abstract": "{In a previous paper we showed that the radio sources selected by combining large areas radio and optical surveys, present a strong deficit of radio emission with respect to 3CR radio-galaxies matched in line emission luminosity. We argued that the prevalence of sources with luminous extended radio structures in high flux limited samples is due to a selection bias. Sources with low radio power form the bulk of the radio-loud AGN population but are still virtually unexplored.  We here analyze their photometric and spectroscopic properties. From the point of view of their emission lines, the majority of the sample are Low Excitation Galaxies (LEG), similar to the 3CR objects at the same level of line luminosity. The hosts of LEG are red, massive ($10.5 \\lesssim {\\rm log} \\, M_*/M_{\\odot} \\lesssim 12$) Early-Type Galaxies (ETG) with large black holes masses ($7.7 \\lesssim {\\rm log} \\, M_{\\rm{BH}}/M_{\\odot} \\lesssim 9$), statistically indistinguishable from the hosts of low redshift 3CR/LEG sources. No genuine radio-loud LEG could be found associated with black holes with a mass substantially lower than $10^8 M_{\\odot}$ or with a late type host. The fraction of galaxies with signs of star formation ($\\sim 5\\%$) is similar to what is found in both the quiescent ETG and 3CR/LEG hosts. We conclude that the deficit in radio emission cannot be ascribed to differences in the properties of their hosts. We argue that instead this could be due to a temporal evolution of the radio luminosity.  A minority ($\\sim 10\\%$) of the sample show rather different properties, being associated with low black hole masses, with spiral galaxies, or showing a high excitation spectrum. In general these outliers are the result of the contamination from Seyfert and from galaxies where the radio emission is powered by star formation. For the objects with high excitation spectra there is no a clear discontinuity in either the host or nuclear properties as they span from radio-quiet and radio-loud AGN. ",
        "introduction": "The advent of large area surveys opened the possibility for the scientific community to explore large samples of extragalactic sources and to set the results on several key issues on strong statistical foundations. In particular, the cross-match of astronomical data from radio and optical surveys provides a unique tool in the analysis of the radio emitting galaxies, since it allows us to identify optically large numbers of radio sources, to obtain spectroscopic redshifts, to determine the properties of their hosts, and to build up their spectral energy distributions.  In recent years several studies have been indeed carried out on large samples of radio galaxies in order to investigate the links between the radio structures, the central engine associated with an Active Galactic Nucleus (AGN), and the host galaxies. In particular, \\citet{best05b} selected a sample of 2215 low-luminosity radio galaxies by cross-correlating SDSS (Data Release 2), NVSS, and FIRST\\footnote{Sloan Digital Sky Survey, \\citep{york00}, National Radio Astronomy Observatory (NRAO) Very Large Array (VLA) Sky Survey \\citep{condon98}, and the Faint Images of the Radio Sky at Twenty centimeters survey \\citep{becker95} respectively.}. The resulting catalogue is by far larger than any previously studied sample of fully optically identified radio-sources.  This sample (hereafter we refer to it as SDSS/NVSS sample) is highly (95 \\%) complete down to the flux threshold of 5 mJy and provides a very good representation of radio-galaxies in the local Universe, up to a redshift of $\\lesssim 0.3$, covering the range $10^{38} - 10^{42}$ erg s$^{-1}$ in radio power.  In a subsequent paper \\citep{best05a} studied the properties of the host galaxies of these radio-loud AGN (hereafter RLAGN) finding that these are massive galaxies, usually of early Hubble type, with stellar masses in the range log $(M_*/M_{\\odot}) \\sim 10^{10}-10^{12}$, and located in richer environments than normal galaxies.  The SDSS/NVSS sample was then considered by \\citet{baldi09} in the context of the properties of low luminosity radio-galaxies. We show that they display a strong deficit of radio emission with respect to their nuclear emission-line luminosity when compared to sources, part of other samples of radio-galaxies, matched in line luminosity.  This result is particularly intriguing considering that the presence of a strong correlation between line and radio luminosity is instead well established (see, e.g., \\citealt{baum89b,rawlings89,morganti97,tadhunter98,willott99}) from the study of different samples of radio-galaxies. Both Low and High Excitation Galaxies (LEG and HEG respectively) obey separately to such a correlation, although the normalizations for the two classes differ \\citep{buttiglione10}. This suggests that the energy source of the narrow lines is closely linked to the source of the radio emission. It can be envisaged that a constant fraction of the available accretion power is channeled into radiative emission (that powers the emission line regions) and into jet kinetic energy (of which the radio emission is the electromagnetic manifestation). The study of the SDSS/NVSS sample suggests instead that total radio luminosity and line emission are independent from each other.  Nonetheless, considering the properties of miniature radio-galaxies \\citep{balmaverde06b} we found that a link exists between radio {\\sl core} and line emission, extending to low luminosities the analogous relation present in high power sources. By assuming that the [O~III] luminosity gives an appropriate bolometric estimate of AGN power, and that radio cores are a good proxy for the jet power, radio galaxies with similar nuclear properties are able to produce an extremely wide range of total radio power, of a factor $\\geq$100. Indeed, these radio-galaxies of very low power are all highly core-dominated, with only feeble extended emission. A similar sample, studied by \\citet{prandoni09} at higher frequency (1.4-15 GHz) and selected at lower flux cutoff ($>1$ mJy), also shows a substantial fraction of very compact radio morphologies (d$<$10 kpc) in the same range of radio power covered by 3CR and B2 sample. We argued that the prevalence of sources with luminous extended radio structures in high flux limited samples is due to a selection bias, since the inclusion of such objects is highly favored. Core dominated sources with a low ratio between radio and emission line power form the bulk of the local radio-loud AGN population but they are still virtually unexplored. Unfortunately, the images available for the SDSS/NVSS sample are not sufficient to properly isolate their radio cores.  The aim of this paper is a better understanding of the physical properties of these RLAGN. In particular we will explore the spectro-photometric properties of their hosts looking for possible differences that might explain their low level of radio emission with respect to the classical powerful radio galaxies matched in line luminosity. Operatively, with the information provided by SDSS we classify the optical spectra of the AGN on the basis of the excitation level. Then, we analyze the key properties of the hosts and the AGN, like the black hole mass, the mass and colors of the hosts, drawing a quantitative comparison with the 3CR sample. The main result is that the sample shows indistinguishable properties from those of 3CR sample. This implies that the reasons for the different radio properties between the bulk of the RLAGN population and the most studied sample of radio-galaxies must be ascribed to other mechanisms.  The paper is organized as follows. In Sect.~\\ref{sample} we briefly present the sample selection carried out by \\citet{best05a}. In Sect.~\\ref{results} and ~\\ref{highz} we analyze the spectro-photometric properties of the SDSS/NVSS sample considering separately the sub-samples in the redshift ranges $0.03<z<0.1$ and $0.1<z<0.3$. What emerges from the study is that the sample is mostly composed of red, massive early-type galaxies with large black hole mass, with spectra typical of LEG. In Sect.~\\ref{comp3c} we compare their broadband properties with those of the 3CR sample, finding clear similarity with LEG. The few apparent exceptions to this general rule are discussed in Sect.~\\ref{exceptions}. The summary and conclusions to our findings are expounded in Sect.~\\ref{summary}. In the Appendix we provide the spectro-photometric diagrams of the sub-sample with $0.1<z<0.3$. ",
        "conclusions": "\\label{summary}  In a previous paper \\citep{baldi09} we showed that the sources selected by \\citet{best05a} by cross-correlating the NVSS, FIRST, and SDSS surveys (the SDSS/NVSS sample) do not follow the correlations between line and radio luminosity defined by radio-galaxies part of samples with a higher flux threshold, such as the 3CR, the B2, and the 2 Jy. They are characterized by a large deficit of radio emission at a given level of emission line luminosity with respect to these more powerful radio-sources. We argued that the prevalence of sources with luminous extended radio structures in high flux limited samples is due to a selection bias. The SDSS/NVSS sources form the bulk of the local radio-loud AGN population but they are still poorly explored.  We here analyze their spectro-photometric properties looking for differences in their hosts and in their optical spectra that might explain their different radio properties. The majority of the SDSS/NVSS sources have emission line ratios characteristic of Low Excitation Galaxies. A comparison with the LEG part of the 3CR sample show that the SDSS/NVSS hosts are indistinguishable from their counterparts of higher radio power from the point of view of their morphology, distribution of black hole and stellar masses, and broad band colors. In fact the hosts of both samples are massive, red, early-type galaxies, with black hole masses in the range $7.7 \\lesssim {\\rm log} (M_{\\rm{BH}}/M_{\\odot}) \\lesssim 9.5$. A small percentage ($\\sim 5\\%$) of objects with signs of recent star formation is also present in both samples, in a similar fraction with respect to the population of quiescent (from the point of view of nuclear activity) galaxies. Thus the differences in radio power between SDSS/NVSS and 3CR sources (matched at a given level of emission line luminosity) cannot be ascribed to the fact that their hosts are different.  Environment can also play a role. In fact, in a dense environment, the confinement of the radio emitting plasma produced by the hot intra-group or intra-cluster medium is more effective than in regions of lower galaxies densities. In this latter case, the rapid expansion of the relativistic plasma causes dramatic adiabatic losses, lowering the observed radio power. However, \\citet{best05a} show that SDSS/NVSS are already in general found in richer environments than normal galaxies. A detailed comparison of the environment of SDSS/NVSS and 3CR sources is needed to address this issue and it will be performed in a forthcoming paper.  In \\citet{baldi09} we argued that SDSS/NVSS sources could be relatively young and their extended radio emission has yet to fully form. This requires that the radio luminosity increases with time as the source expands, slowly burrowing their way into the interstellar medium of the host galaxy. The results presented here provide us with a further piece of evidence in favor of a evolutionary sequence in radio-loud AGN. In fact we found that Unclassified Galaxies are more often associated with extended radio sources with respect to Low Excitation Galaxies (the percentages being 27 \\% and 5 \\% for UG and LEG, respectively). This is counter-intuitive, since LEG are brighter in line emission than UG; one might have expected a higher line luminosity, generally indicative of a higher AGN power, to correspond more often to a well developed radio structure. However, also the line luminosity evolves with time and with the size of the radio source. The compression of the ambient medium due to the passage of the jet head (and to the lateral expansion shocks) increases the density of the line emitting clouds and consequently (given the dependence of line luminosity on the density squared) their line luminosity. The line luminosity increases with time, as the jets compress more clouds (see e.g. \\citealt{labiano08}) until the radio source is confined to within the host galaxy. Even larger sources propagate in regions of lower gas densities. At a given level of the ionizing continuum, they are expected to be less luminous in line that smaller sources where the jet-cloud interactions are more efficient. Thus a radio-galaxy moves with time in the $L_{1.4~{\\rm GHz}}-L_{\\rm [O~III]}$ plane along a complex evolutionary track. The higher fraction of UG with extended radio structures can be due to the evolution of part of the LEG population, from small radio sources of high $L_{\\rm{[O~III]}}$, to larger size and smaller $L_{\\rm{[O~III]}}$.  Additional indicators of the jet and AGN power are needed to further compare high and low radio power sources. These can be obtained, for example, from measurements of the radio cores (from images of higher resolution and at higher frequency than the FIRST maps currently available) or of the higher frequency nuclear emission by using, e.g., HST or Chandra imaging. In particular it is crucial to test whether also the sources belonging to the SDSS/NVSS sample obey to the relation between radio core and line emission spanning from miniature to high power radio-galaxies \\citep{baldi09}. This would represent a clear indication that there is indeed a close connection between the radiative and kinetic output in radio-loud AGN, but that the best indicator of jet power is the radio core luminosity and not the total radio emission. This interpretation is in line with the proportionality of core and jet power derived from the estimates of the energy required to form the cavities observed in X-ray images \\citep{heinz07}. Better radio images will also provide us with a measurement of the core dominance, to test the prediction that the deficit of radio power of SDSS/NVSS sources is indeed related to a higher core dominance with respect to, e.g., 3CR objects, in analogy with the results we obtained for miniature radio-galaxies.  After a detailed analysis on a object by object basis, we concluded that no bona-fide radio-loud AGN with low excitation spectra could be found in a late type host or with black holes with a mass substantially lower than $10^8 M_{\\odot}$. The SMBH mass distributions of 3CR/LEG and in the SDSS/NVSS samples suggests that a RLAGN LEG can only be associated with relatively large mass of the central object. This might represent a very valuable information in our quest for the origin and evolution of RLAGN (see e.g. \\citealt{chiaberge08}) and of their host galaxies. Nonetheless, since a link between $M_{\\rm{BH}}$ and the normalization of the radio luminosity function exists, RLAGN might be associated with low SMBH masses and not be represented in the sample because they are too rare. However, we estimated that 7 low $M_{\\rm{BH}}$ objects would be expected, while none was found. This idea can be tested in greater depth by exploring the radio properties of lower mass ETG, harboring black holes around $10^7 M_{\\odot}$.  A minority ($\\sim 10\\%$) of the SDSS/NVSS sample shows rather different properties, showing a spectrum typical of high excitation or star forming galaxies. They strikingly differ from LEG: they have a wider distribution of SMBH masses, ranging from $\\sim 10^6 M_{\\odot}$ to $\\sim 10^9 M_{\\odot}$ and they are often found in spiral galaxies. Only a few of the HEG in the SDSS/NVSS sample (less than $\\sim$ 10 objects) have spectro-photometric properties similar to the HEG found in the 3CR. Furthermore, there might be a contamination from radio-quiet galaxies with high excitation spectra whose radio emission is dominated by star formation.  Interestingly, however, in the SDSS/NVSS there is no clear discontinuity in either the host or the AGN properties moving from galaxies lying in the region characteristic of Seyfert to that covered by 3CR/HEG, objects that would be classically defined as radio-quiet and radio-loud AGN. This issue certainly deserves further studies, requiring to collect additional multiband data on these galaxies.  It is possible that the same situation occurs also for LEG. The density of SDSS/NVSS sources, at given line luminosity, increases toward lower radio luminosity (with a dependence well described by a power law $N \\sim L_{1.4~{\\rm GHz}}^{-1}$) down to the completeness limit imposed by the radio flux limit. We cannot establish whether a discontinuity exists between the RLAGN discussed here and the population of radio-quiet LEG, i.e. the LINER. An extension of this study toward even lower radio flux limits, probably at sub-mJy level, is required to evidence the presence of a possible radio-quiet/radio-loud dichotomy in low excitation AGN.  Furthermore, if indeed galaxies similar from the point of view of their nuclear properties give rise to radio-sources of widely different power and morphology, the use of the total radio power as calorimeter of the jet is questioned. This opens also the issue of how to define the radio-loudness of AGN. Apparently, RL and RQ AGN are best separated considering their nuclear properties, in particular by comparing the luminosity in the radio with X-ray and optical data (e.g. \\citealt{early1}), stressing again the importance of high resolution multi-band measurements of their nuclear emission.  Finally, we note that an important issue in the study of these radio sources of, generally, low power, is related to the identification of the dominant process of the radio emission, i.e. AGN or star formation. \\citet{best05a} adopted a separation based on the radio luminosity expected for a stellar population of a given age (derived from the value of the $D_n(4000)$ index) and mass. While the proposed method is sufficiently accurate to build, e.g., radio luminosity functions, it cannot be safely applied to individual galaxies. In fact, we found that many objects located just above the threshold in the $D_{n}(4000)$ vs. $L_{\\rm 1.4GHz}/M_{*}$ plane have spectro-photometric characteristics very different from the bulk of the sample. This suggests that the ``grey region'' where the two populations (with radio powered by an AGN or by star formation) overlap covers a rather large range of these parameters. This is likely due to the different scales probed by NVSS images and SDSS spectra. For example, we found a spiral galaxy whose SDSS spectrum is indicative of an old stellar population. However, the SDSS fiber covers only its bulge and, consequently, the predicted radio emission from star formation is relatively low, suggesting an AGN origin. However, no radio core is seen in the FIRST images, suggesting that the radio emission might originate from young stars in the spiral arms, not covered by the nuclear spectrum. Deeper radio images and with better resolution than the FIRST and NVSS maps could be used to isolate these objects based on their radio morphology and on the spatial association of radio and optical structures."
    },
    "1005/1005.3479_arXiv.txt": {
        "abstract": "We report the discovery of X-ray eclipses in the recently discovered accreting millisecond X-ray pulsar Swift J1749.4$-$2807. This is the first detection of X-ray eclipses in a system of this type and should enable a precise neutron star mass measurement once the companion star is identified and studied.  We present a combined pulse and eclipse timing solution that enables tight constraints on the orbital parameters and inclination and shows that the companion mass is in the range $0.6 - 0.8 M_{\\odot}$ for a likely range of neutron star masses, and that it is larger than a main sequence star of the same mass. We observed two individual eclipse egresses and a single ingress. Our timing model shows that the eclipse features are symmetric about the time of $90^{\\circ}$ longitude from the ascending node, as expected. Our eclipse timing solution gives an eclipse duration (from the mid-points of ingress to egress) of $2172\\pm13$ s. This represents 6.85\\% of the 8.82 hr orbital period.  This system also presents a potential measurement of ``Shapiro'' delay due to General Relativity; through this technique alone, we set an upper limit to the companion mass of 2.2 $M_\\sun$. ",
        "introduction": "The neutron star mass (and radius) distribution has important implications for our understanding of the equation of state (EOS) of ultra-dense matter. The EOS sets directly a number of potentially observable quantities, including the maximum mass and spin frequency of a neutron star (Lattimer \\& Prakash 2001). Moreover, measurements of neutron star radii would strongly constrain the nuclear symmetry energy (Steiner \\etal\\ 2010).  At present there remain serious gaps in our understanding of the mass distribution of neutron stars, particularly as regards their upper mass limit. Where neutron star masses are known with precision it is in most cases for relatively young, binary radio pulsars in which one or more relativistic, Post Keplerian (PK) parameters can be directly measured (Thorsett \\& Chakrabarty 1999). Two famous examples of such systems include the original Hulse - Taylor binary pulsar PSR 1913+16 (Taylor \\& Weisberg 1989), and the double pulsar PSR J0737$-$3039 (Kramer et al.  2006). These neutron stars have masses that appear to cluster near or below the ``canonical'' value of 1.4 $M_{\\odot}$, and hence do not strongly constrain the maximum mass nor the EOS.  Studies of older, accreting neutron star systems hold significant promise for directly probing the high end of the mass distribution because it is likely these objects have accreted significant mass during their evolutionary history. Pulsar timing studies of binary millisecond radio pulsars have recently led to strong indications for higher mass neutron stars. For example, Freire et al. (2009) present timing measurements of PSR J1903+0327 that indicate its neutron star has a mass of $1.67\\pm0.01 M_{\\odot}$ (also Champion et al. 2008). A related population of sources which could provide important mass constraints are accreting millisecond X-ray pulsars (AMXPs). Additionally, these systems also produce X-ray emission directly from the neutron star surface in the form of thermonuclear X-ray bursts, and in principle can also be used to constrain neutron star radii, something that cannot be done for the radio pulsar systems unless spin orbit coupling can be directly observed (Lattimer \\& Schutz 2005). Currently twelve of these systems are known, with the first, \\object[SAX J1808.4-3658]{SAX J1808.4$-$3658}, having been discovered in 1998 (Wijnands \\& van der Klis 1998).  A subset of the AMXPs are ultracompact binaries, with orbital periods less than an hour and have very low mass ($<< 0.1 M_{\\odot}$), degenerate dwarf companions.  An example of a system of this type is \\object[XTE J1751-305]{XTE J1751$-$305} (Markwardt et al. 2002). The remainder appear to more closely resemble the ``classical'' accreting, neutron star low mass X-ray binaries (LMXBs), with orbital periods of ~hours and low mass main sequence companions (such as SAX J1808.4$-$3658).  Several factors currently limit the use of these systems to obtain precise neutron star mass estimates. First, the population is largely associated with the Galactic bulge. Therefore the resulting high optical extinction makes detection of companions challenging. Second, their binary inclinations have, to date, not been strongly constrained, as would be possible, for example, with the detection and timing of X-ray eclipses. Indeed, observations of eclipsing AMXPs would enable the binary inclination to be tightly constrained and, combined with precision X-ray pulse timing and spectroscopic measurements of the companion, could lead to precise mass constraints for their neutron stars. Finally, it has not yet been possible to measure any PK parameters (such as the Shapiro delay) in these systems.  In this Letter we report the first detection of X-ray eclipses from an AMXP, in the system \\object[SWIFT J1749.4-2807]{SWIFT J1749.4$-$2807} (hereafter J1749).  J1749 was first discovered in 2006 based on the detection with \\swift/BAT of a thermonuclear X-ray burst (Schady et al. 2006). Follow-up observations with \\swift\\ set an upper limit to the source distance of $6.7\\pm1.3$ kpc (Wijnands et al. 2009). The source was detected again in 2010 by {\\it INTEGRAL}, exhibiting both persistent and bursting activity (Pavan et al. 2010; Chenevez et al. 2010). Altamirano et al. (2010a), using \\rxte\\ observations during the recent outburst of this object, found that it harbors a 518 Hz accreting millisecond pulsar, with strong power also evident at the 1036 Hz first harmonic (Bozzo et al. 2010).  Analysis of additional \\rxte\\ observations first by Belloni et al. (2010) and then by Strohmayer \\& Markwardt (2010) found an 8.82 hr circular orbit, and eclipses centered at the orbital phase of superior conjunction of the neutron star (Markwardt et al. 2010). Here we present a timing analysis of the pulsar and its eclipses, and derive constraints on the properties of the binary components. Our goal in this work is to establish the orbital solution from pulse timing to a sufficient degree of precision to place the eclipse observations in context.  We leave a detailed characterization of the pulse profiles and timing noise to other work (Altamirano et al. 2010b; Ferrigno et al. 2010). ",
        "conclusions": "The simultaneous presence of X-ray pulsations and X-ray eclipses have allowed us to constrain the companion of the system with unprecedented accuracy.  The companion mass of $\\sim$0.7 $M_\\sun$ is rather large compared to most of the other AMXP companions.  In contrast to the ``ultra-compact'' systems such as XTE J1751$-$305, whose companion mass is a reasonable multiple of Jupiter's mass (Markwardt et al. 2002), J1749 has a mass more comparable to the Sun.  The companion of J1749 appears to be larger than a typical main sequence star, as obtained from Zombeck (1982; see Figure \\ref{fig:orbgrid}), but also from more modern studies (Chabrier et al. 2007).  For a neutron star of 1.4 $M_\\sun$, the companion is about 20\\% larger than expected from the main sequence track.  We can interpret this as ``bloating'' of the companion due to X-ray irradiation by the neutron star which will cause the outer layers of the star to expand (Bildsten \\& Chakrabarty 2001).  The eclipse timing presented here places relatively tight constraints on the structure of the mass donor.  Theoretical calculations indicate that the donor in a mass transfer binary can be perturbed away from a main sequence spectral type -- mass relationship (see Beuermann et al. 1998; Kolb, King \\& Baraffe 2001).  Beuermann et al.  (1998) compare observational results for the donors in CVs with results of theoretical calculations, and show that systems with orbital periods longer than about 6 hours span a range of spectral type, whereas short period systems (less than 2 hr) behave more like solar-abundance, main-sequence stars.  In related work, Kolb, King \\& Baraffe (2001) show that for a given measured spectral type a range of masses is possible based on the evolutionary status of the donor and its mass transfer history.  The limits from their Table 1, when compared with our companion mass estimates, would suggest a spectral type for the secondary of J1749 from about K1 to K3.  This range is also consistent with the results of Beuermann et al. (1998) for an $\\approx 9$ hr binary period (see their Fig. 5).  There is only one orbital variable left in the system to be known.  If the orbital inclination or companion mass or radius could be determined, then the neutron star mass could also be determined. In general, spectroscopic measurements of the companion to determine its radial velocity profile would provide the mass ratio and thus fully constrain the system components, including the neutron star mass.  As noted in the introduction, systems like J1749 provide a unique opportunity to probe the high mass end of the neutron star mass distribution.  Before this can be accomplished, an IR/optical counterpart must be identified.  Unfortunately, high resolution imaging {\\it Chandra} observations performed shortly after the current outburst failed to detect an X-ray source (Chakrabarty et al. 2010), so IR/optical observations will be difficult.  Nevertheless, spectroscopy should be a focus of future observations when the source reappears.  We believe this is the first time that realistic limits have been set for the Shapiro delay effect at X-ray wavelengths, for a system outside of our solar system.  While our companion mass upper limit was ultimately unconstraining compared to the eclipse mass constraint, the technique does show the power to measure relativistic effects with a mission like {\\it RXTE}.  If J1749 has another outburst (especially a longer duration outburst which allows better coverage near eclipse), it may be possible to place stronger constraints on the system masses."
    },
    "1005/1005.3365_arXiv.txt": {
        "abstract": "{ We use a long (300~ksec), continuous {\\it Suzaku} X-ray observation of the active nucleus in NGC~1365 to investigate the structure of the circumnuclear BLR clouds through their occultation of the X-ray source. The variations of the absorbing column density and of the covering factor indicate that the clouds surrounding the black hole are far from having a spherical geometry (as sometimes assumed), instead they have a strongly elongated and cometary shape, with a dense head ($\\rm n\\sim 10^{11}~cm^{-3}$) and an expanding, dissolving tail. We infer that the cometary tails must be longer than a few times $\\rm 10^{13}~cm$ and their opening angle must be smaller than a few degrees. We suggest that the cometary shape may be a common feature of BLR clouds in general, but which has been difficult to recognize observationally so far.  The cometary shape may originate from shocks and hydrodynamical instabilities generated by the supersonic motion of the BLR clouds into the intracloud medium. As a consequence of the mass loss into their tail, we infer that the BLR clouds probably have a lifetime of only a few months, implying that they must be continuously replenished. We also find a large, puzzling discrepancy (two orders of magnitude) between the mass of the BLR inferred from the properties of the absorbing clouds and the mass of the BLR inferred from photoionization models; we discuss the possible solutions to this discrepancy. } ",
        "introduction": "\\label{intro}  The variability of the absorbing gaseous column density $\\rm N_H$ observed in the X-ray spectra of Active Galactic Nuclei have revealed that a significant fraction of the absorbing medium must be clumpy \\citep{risaliti02}. Moreover, $\\rm N_H$ variations on time scales as short as a few days or hours have shown that a significant fraction of such absorbing clouds must be located very close to the nuclear X-ray source \\citep{elvis04,risaliti05a,puccetti07} and, more specifically, within the Broad Line Region (BLR).  The nucleus of the galaxy NGC~1365 is amongst the sources that were investigated more thoroughly in this respect, thanks to its brightness and probably also to a fortunate (intermediate) inclination of the absorbing medium relative to our line of sight. This is a nearby (z=0.0055) type 1.8 Seyfert galaxy that was observed several times in the X-rays, and displayed strong column density variations on time scales as short as a few days \\citep[e.g.][]{risaliti05a,risaliti07}. Most of the observations had probed absorption variations of this source in discrete, non-contiguous time intervals. More recently, (nearly) continuous XMM observations of NGC~1365 were presented in \\cite{risaliti09a} and in \\cite{risaliti09b}. In particular, the latter observation allowed the continuous monitoring of an absorption event lasting about 40~ksec, which was modelled with a cloud partially eclipsing the X-ray source. More specifically, the spectral variations were modelled by assuming a constant column density of the absorber ($\\rm N_H\\sim 3.5~10^{23}~cm^{-2}$) and a variable covering factor. The inferred density, velocity and distance from the black hole of the variable absorber is strikingly similar to those of the BLR clouds. The complex spectrum always requires a foreground non-variable (or very slowly variable) absorber ($\\rm N_H\\sim$ a few times $10^{22}~cm^{-2}$), totally covering the source, possibly associated with gas at larger distances from the nucleus (possibly in the host galaxy).  NGC~1365 has also been observed with {\\it Suzaku} \\citep{mitsuda07} with a continuous observation lasting more than 300~ksec. The integrated spectrum was already presented in \\cite{risaliti09c}, mostly focusing on the high energy part obtained with the Hard X-ray Detector (HXD), and revealing an unusually strong excess of emission at $\\rm E>10~keV$. Such excess is ascribed to X-ray radiation piercing a partial, Compton thick ($\\rm N_H \\sim 4\\times 10^{24}~cm^{-2}$) absorber. The high energy excess appears stable over long time scales \\citep{risaliti00,risaliti09c}. The emerging scenario is that three distinct absorbers are present: (1) a distant absorber with $\\rm N_H < 10^{23}~cm^{-2}$, probably associated with gas in the host galaxy; (2) an absorber made of broad line region (BLR) clouds with $\\rm N_H\\sim 10^{23-24} cm^{-2}$ rapidly orbiting around the black hole; (3) an absorber with $\\rm N_H \\sim 4\\times 10^{24}~cm^{-2}$ partially covering the source, consisting either of the outer region of a warped accretion disk, or of a large number of small ($\\rm <10^{12}cm$) and dense ($\\rm n>10^{12}cm^{-3}$) clouds located in the vicinity of the accretion disk.  In this paper we focus on the absorption variations in the spectral range 2-10~keV, ascribed to BLR clouds passing along the line of sight, through a detailed temporal analysis of the spectra provided by the {\\it Suzaku} low energy (XIS) data of NGC~1365. We identify two eclipsing events that we can model in detail in terms of variations of both covering factor and column density. The time variations of these quantities are highly asymmetric and strongly suggest that the absorbing clouds have a cometary shape, i.e. a well defined head and a more diffuse elongated, conical, tail. ",
        "conclusions": "We have presented a detailed spectral analysis of a long ($\\rm \\sim 300~ksec$), continuous {\\it Suzaku} observation of the Seyfert nucleus in the galaxy NGC~1365. The spectrum shows evidence for an absorption component that is variable in time, both in terms of column density and in terms of covering factor, which is ascribed to clouds eclipsing the X-ray source. This is the first time that a temporally resolved X-ray spectral analysis is able to break the degeneracy between the evolution of the column density ($\\rm N_H$) and the covering factor (CF) of the X-ray absorber.  We identify two main eclipses. The temporal evolution of each eclipse is far from being symmetrical. The initial occultation is very rapid (within $\\rm \\sim 1~ksec$) and covering about 65\\% of the source. Subsequently the covering factor increases, but more slowly, reaching unity in about 50~ksec. The absorbing column density of the cloud is about $\\rm 10^{23}~cm^{-2}$ at the beginning of each eclipse and decreases afterwards. These results are inconsistent with a spherical geometry for the absorbing clouds. The most likely geometry compatible with the observations is an elongated, ``cometary'' shape, with a dense ``head'' ($\\rm n\\sim 10^{11}~cm^{-3}$, consistent with that expected for the clouds of the BLR) and with a dissolving and expanding tail.  The data allow us to quantitatively constrain the geometry, dynamics and location of such cometary clouds. The cometary clouds are probably located at a distance of about $\\rm 2~10^{15}~cm$ from the nuclear black hole, well within the estimated BLR radius ($\\rm \\sim 10^{16}~cm$), strongly supporting the association of these absorbing systems with BLR clouds (at least with the inner, high-ionization ones). The cometary clouds ``head'' must have a size comparable to the X-ray source and must be moving with velocity higher than about 1000~km/s (consistent with the velocity expected for the BLR clouds). The cometary tail must be longer than a few times $\\rm 10^{13}~cm$. The tail opening angle must be very narrow, less than a few degrees, and consistent with the opening angle of the Mach cone expected from the supersonic motion of the cloud into the hot intracloud medium.  We suggest that such cometary clouds may be common to most AGNs, but have been difficult to recognize in previous X-ray observations.  We estimate that the cloud ``head'' loses a significant fraction of its mass through the cometary tail, which is expected to cause the total cloud destruction within a few months. If these clouds are representative of most BLR clouds (or at least the high-ionization ones), our result implies that the BLR region must be continuously replenished with gas clouds, possibly from the accretion disk.  We briefly discussed the possible nature of such cometary clouds. The most likely scenario is that the tail is made of gas lost by the cloud head through hydrodynamical instabilities generated by its supersonic motion through the hot intracloud medium.  We can estimate the mass of the absorbing clouds ($\\rm M_{cloud}\\sim 10^{-10}~M_{\\odot}$) and their total number within the central region ($\\rm \\mathcal{N_C}\\sim 3~10^7$). The inferred total mass of the BLR is about $\\rm 4~10^{-3}~M_{\\odot}$, which is two orders of magnitude lower than the BLR mass inferred from photoionization models. The discrepancy may originate from a population of large, massive BLR clouds not identified in our eclipsing studies. Alternatively, photoionization models may overestimate the BLR mass. In particular, UV and X-ray radiation produced by the shocks generated by the supersonic motion of the clouds may provide a local source of ionizing photons, not accounted for by classical photoionization models that assume a central point-like radiation source."
    },
    "1005/1005.5363_arXiv.txt": {
        "abstract": "A recent observation of the nearby (z=0.042) narrow-line Seyfert 1 galaxy \\src\\ on 2007 May 31 showed strong quasi-periodic oscillations (QPOs) in the 0.3--10 keV X-ray flux. We present phase-resolved spectroscopy of this observation, using data obtained by the EPIC PN detector onboard \\xmm. The ``low'' phase spectrum, associated with the troughs in the light curve, shows (at $>4\\sigma$ confidence level) an absorption edge at $0.86\\pm0.05$ keV with an absorption depth of $ 0.3\\pm0.1$.  Ionized oxygen edges are hallmarks of X-ray warm absorbers in Seyfert active galactic nuclei (AGN); the observed edge is consistent with H-like O VIII and implies a column density of $N_{\\rm OVIII}\\sim3\\times10^{18}$ cm$^{-2}$. The edge is not seen in the ``high'' phase spectrum associated with the crests in the light curve, suggesting the presence of a warm absorber in the immediate vicinity of the supermassive black hole which periodically obscures the continuum emission.  If the QPO arises due to Keplerian orbital motion around the central black hole, the periodic appearance of the O VIII edge would imply a radius of $\\sim9.4(M/[4\\times10^6 \\msun])^{-2/3}(P/[1 hr])^{2/3}\\;r_{\\rm g}$ for the size of the warm absorber. ",
        "introduction": "\\src\\ is a narrow line Seyfert 1 galaxy (NLS1; see \\citealt{op1985} for definition of NLS1s). It is well known for its high soft X-ray excess compared not only to other AGNs but also among the NLS1s \\citep{pu1995,boller1996,middleton2007}.  While the origin of this soft excess is still not clear, the similarity of its soft X-ray spectrum with that of stellar black hole binaries in their `high state' has been used to postulate that \\src\\ harbors a comparatively low mass black hole which is currently in a state of high mass accretion rate \\citep{pounds1995}.  A recent X-ray observation of \\src\\ using the \\xmm\\ satellite made between 2007 May 31 and 2007 June 1 showed a strong signature of quasi-periodic variability in the 0.3--10 keV light curve \\citep{g08}. The oscillations were transient in nature since they have never been seen in any prior observations of this source, nor did we see strong QPO-like variability in the 0.3--10 keV light curve obtained during a subsequent \\xmm\\ observation made on 2009 May 31.  It was noted by \\citet{g08} that the frequency of the observed oscillations was similar to what would be expected if the frequencies of the quasi-periodic oscillations often seen in Galactic black hole binaries \\citep[see e.g.][for a review on observations of QPOs in compact stellar X-ray binary systems]{vdk2006} were to be scaled by the same factor as the mass ratio between the mass of the supermassive black hole in \\src\\ and the mass of typical stellar black holes.  Recently \\citet{MiddletonDone2009} have performed detailed comparison of the timing properties of this \\xmm\\ observation with that of stellar mass black hole binaries and concluded that the observed QPO in \\src\\ is similar to the 67 Hz QPO seen in GRS~1915+105 (a black hole binary system which, like \\src, also boasts a high L/L$_{\\rm Edd}$). This scaling of oscillation frequencies with mass would support the hypothesis that accretion physics scales with mass and that the supermassive black holes at the center of galaxies are scaled-up versions of stellar-mass black holes.  The origin of QPOs in stellar black hole (and neutron star) binaries is not yet well understood.  Various models have been proposed, but none can fully explain the wealth of observed phenomenology.  QPO frequencies in stellar black hole binaries range from Hertz to kHz, and it is extremely difficult to carry out phase-resolved spectroscopy in these systems. In the sparse cases \\citep[e.g. in the source GRS~1915+105, reported by][]{Morganetal1997} where it is possible to ``see'' the QPO in the light curves, the phase of the oscillations appear to perform a random walk. \\citet{MillerHoman2005} showed that for GRS~1915+105 the flux in the Fe K$\\alpha$ line (created by hard X-ray photons irradiating a relatively cold accretion disk) varies with the QPO oscillation phase of the 1 and 2 Hz QPOs, thus suggesting that the 1 Hz and 2 Hz QPOs may be linked to variable reflection.  The strong variability amplitude and comparatively large period of the oscillations in \\src\\ make it an ideal source to carry out phase-resolved spectroscopy and probe changes in the spectral energy distribution between different phases.  The origin of the soft excess in \\src\\ is not clear, and its time-averaged spectral energy distribution can be well modeled by a wide variety of models including reflection from a partially ionized accretion disc \\citep{crummy2006}, Comptonized disc emission from a low temperature disc, ionized partial covering, or a smeared disc wind seen in absorption \\citep[see e.g.][]{middletonetal2009}.  Phase-resolved spectra could give important insights for breaking this degeneracy.  Therefore one of our major goals in this work is to search for signatures of variable reflection (in the continuum and reflection lines) and/or variable absorption features (e.g.  changes in the properties of O VII, O VIII edges which are the hallmark of the presence of warm absorber in active galactic nuclei).  \\citet{middletonetal2009} analyzed the energy dependence of the observed variability in the EPIC data and showed that the variability primarily originated in the high-energy photons.  They also presented a spectral decomposition of the EPIC MOS spectrum, averaged over the entire observation.  Here we analyze the EPIC PN data taken during this observation to test if there is any signature of spectral variations between the different phases.  Due to limitations in photon statistics (largely caused by the high pileup on the CCD as described in \\S\\ref{prepare_data}; also see \\citealt{middletonetal2009}) we could extract spectra with good signal-to-noise ratio from two phases only: the ``high'' phase spectrum by accumulating spectra near the crests, and the ``low'' phase spectrum by accumulating spectra near the troughs. In \\S\\ref{prepare_data} we describe in detail the algorithm we used to create the phase-resolved ``high'' and ``low'' spectra.  The spectral analysis is described in \\S\\ref{analysis}, and conclusions are summarized in \\S\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion} In this work we have presented phase-resolved spectroscopy of the narrow line Seyfert 1 galaxy \\src. We have used the \\xmm\\ observation of this source between 2007 May 31 and 2007 June 1, when the X-ray light curve showed strong quasi-periodic oscillations.  We have extracted spectra during the high (crest) and low (trough) phases.  Simple continuum models assuming either a blackbody plus broken power-law, or two blackbodies plus a power-law model, fit the high phase spectrum quite well.  The required low-energy power-law index is very steep ($\\Gamma\\sim3.6\\pm0.1$) for a single blackbody plus broken power-law model, which is typical for narrow line AGNs with a high soft X-ray excess \\citep{boller1996,pounds1995,middleton2007}.  The spectrum can also be well modeled using two blackbodies and a power-law continuum.  Due to the steep decline of the spectrum, and high pileup on the PN detector which severely limits the photon statistics, we do not have good quality spectral data beyond $\\sim$2--3 keV.  Therefore for these spectra we cannot very well constrain the relatively flatter 2--10 keV continuum usually seen in this source \\citep{crummy2006,middletonetal2009}, although fits to a two-blackbody plus power-law model show that the $>2$ keV spectrum is comparatively flatter with a photon index of $\\sim2$.  The low phase spectrum cannot be adequately described by scaling the overall normalization of either of the best fit high phase models, suggeting that the shape of the low phase spectrum is different from that of the high phase.  Analyzing the r.m.s. variability at different energies \\citet{middletonetal2009} also concluded that the variability is mainly from the hard photons.  The right panel of Fig.~\\ref{f:lo_renorm} shows that the differences between the low and high phases are most prominent at higher energies.  In particular, the residuals in Fig.~\\ref{f:lo_renorm} show a sharp drop in the continuum flux between 0.8--0.9 keV, which then slowly increases toward the expected continuum flux at higher energies, reminiscent of an absorption edge.  The inclusion of an edge improves the fits significantly (at $>$4$\\sigma$ level) for both single-blackbody+broken power-law model and double-blackbody+power-law model.  The best fit values for the edge threshold energy and absorption depth are $0.87\\pm0.04$ keV and $0.29\\pm0.11$ respectively for the single-blackbody+broken power-law model, and $0.86^{+0.05}_{-0.06}$ keV and $0.23^{+0.12}_{-0.11}$ respectively for the double-blackbody+power-law model.  The edge is most likely associated to the 0.87 keV (rest frame) K-edge of (H-like) oxygen VIII which is quite commonly seen in Seyferts \\citep[see e.g.,][]{reynolds1997}, and is indicative of the presence of optically thin, photoionized matter along the line of sight.  \\src\\ however is at a redshift of 0.042; therefore the 0.87 keV edge (in rest frame) should appear at 0.84 keV (easily within the 90\\% confidence error bar for both models). The presence of a warm absorber should create other features in the spectrum, e.g. the 0.74 keV K edge from He-like O VII. While the residuals in Fig.~\\ref{fig:lo_edge} do show a drop in the flux, it is not very significant statistically.  To estimate the column density ($N_{\\rm OVIII}$) of the O VIII ions responsible for the edge we used Eq.(1) of \\citet{verneretal1996} and calculated the absorption cross section $\\sigma$ (for O VIII, $\\sigma[{\\rm E}=0.87\\;{\\rm keV}]=0.098$ Mb).  Since $\\tau=N_{\\rm OVIII}\\cdot\\sigma$, we obtain $N_{\\rm OVIII}=3\\times10^{18}$ cm$^{-2}$.  Estimating the number density ($n_{\\rm OVIII}$) requires a size-scale for the warm absorber.  Assuming that the QPO originates in the accretion disk, the warm absorber is could be confined within the orbit of the QPO, or in case of an outflow it originates inside the QPO orbit. The mass of the supermassive black hole in \\src\\ is in the range of $1-7\\times10^{6}$ \\msun\\ \\citep{BianHuang2010,zhouetal2010}.  A circular Keplerian orbit would have a radius of $9.4(M/[4\\times10^6 \\msun])^{-2/3}(P/[1 hr])^{2/3}\\; r_{\\rm g}$, which would be an upper limit on the size of the absorber in this case. This would give $n_{\\rm OVIII}=3\\times10^5$ cm$^{-3}$ for the number density of OVIII ions.  Note however that this assumes a Keplerian orbit origin for the QPO, and therefore depends on the mass of the central supermassive black hole.  Thus in \\src\\ we may be seeing an infalling ``blob'', the emission from which is periodically absorbed as it passes behind a warm absorber in the immediate vicinity of the central engine.  If the ``blob'' is at $\\sim$9 $r_{\\rm g}$, that would imply the ionizing flux must be generated within 9 $r_{\\rm g}$ (or else we would not see absorption), and this would be easier to accomplish if the black hole is spinning since the disk can get down to a smaller radius (e.g. as low as 1.25 $r_{\\rm g}$ for a maximally spinning black hole). Recent work by \\citet{crummy2006} has shown that NLS1 spectra can be modeled using reflection from a disk around a spinning black holes \\citep[see][for a review]{Miller2007}.  We also created a multidimensional grid of models using the XSTAR photoionization code, to create a more physical model for the warm absorber and explore a wide range of column density, ionization parameter, oxygen and neon abundance. Fits to the low-phase data using this model suggests a warm absorber column density of $\\sim4\\times10^{21}$ cm$^{-2}$, log($\\xi$)=$3.2 _{-0.1}^{+0.2}$, and an oxygen abundance of $2.1 _{-0.5}^{+0.8}$ relative to solar.  The neon abundance could not be well constrained from the data, although the best-fit solution suggests a sub-solar value.  Another possible physical scenario for the difference between high and low phase spectrum could be that the warm absorber is along the line of sight, between us and the source (where the QPO originates), and the warm absorber responds to flux variations in the source itself, i.e. it is more ionized and transparent at high fluxes and less at lower fluxes. However, based on our modeling of the spectra using physically motivated models generated using XSTAR this scenario appears unlikely because the change in luminosity of the ionizing flux between low and high phases is too small to create the observed difference in the depth of the OVIII edge (also see \\S\\ref{analysis}).  It is interesting to note in this context that a somewhat complementary situation has been observed in the Seyfert galaxy NGC~1365, where \\citet{Risalitietal2009} have reported a possible transit of an obscuring cloud (with $N_H\\sim3.5\\times10^{23}$ cm$^{-2}$ and other inferred properties similar to that of a broad-line region cloud) in front of the central X-ray source.  Signatures of variable absorption on short timescales have been previously observed in high resolution {\\em Chandra} spectra of the stellar black hole candidate H1743--22 \\citep{miller2006}.  For H1743--22, the lines were slightly blueshifted, suggesting an outflow origin. For \\src, we do not have definitive evidence for an outflow, but the best fit edge energy for both models is higher (though within 90\\% confidence interval) than the rest frame energy of the oxygen edge.  If the warm absorber in \\src\\ is indeed in an outflow, then the observed oscillations could originate due to an instability in the inner accretion disk threaded by a poloidal magnetic field, giving rise to an outflowing wind or jet \\citep{BlandfordPayne1982,tagger1999}.  The observed oscillations in this case would correspond to scaled up versions of low-frequency QPOs \\citep[see e.g.,][]{vdk2006} seen in stellar mass X-ray binaries.  Since the presence of a warm absorber is bound to have noticeable imprints on other regions of the electromagnetic spectrum, especially in UV and soft X-rays, future deep UV and X-ray observations of \\src\\ will be useful not only in constraining the origin of the oscillations, but also in gaining a better understanding of the physical emission mechanism of sources like \\src\\ which show a strong soft-X-ray excess. Future X-ray missions with improved sensitivity and larger collecting area like the {\\em International X-ray Observatory} \\citep[IXO;][]{Milleretal2009,Whiteetal2010} would play key role in understanding accretion geometry as well as emission/absorption mechanism in these sources."
    },
    "1005/1005.0293_arXiv.txt": {
        "abstract": "Stellar activity has a particularly strong influence on planets at small orbital distances, such as close-in exoplanets. For such planets, we present two extreme cases of stellar variability, namely stellar coronal mass ejections and stellar wind, which both result in the planetary environment being variable on a timescale of billions of years. For both cases, direct interaction of the streaming plasma with the planetary atmosphere would entail servere consequences. In certain cases, however, the planetary atmosphere can be effectively shielded by a strong planetary magnetic field. The efficiency of this shielding is determined by the planetary magnetic dipole moment, which is difficult to constrain by either models or observations. We present different factors which influence the strength of the planetary magnetic dipole moment. Implications are discussed, including nonthermal atmospheric loss, atmospheric biomarkers, and planetary habitability. ",
        "introduction": "One of the many fascinating questions in the field of exoplanet studies is the search for habitable worlds. Because of their relatively small mass, low luminosity, long lifetime and large abundance in the galaxy, M dwarfs are sometimes suggested as prime targets in searches for terrestrial habitable planets \\citep{Tarter07, Scalo07}. Interestingly, M dwarfs also seem to have a larger number of Super-Earth planets (i.e.~planets with a mass smaller than 10 terrestrial masses, i.e.~$M \\lesssim 10 M\\E$) than more massive stars \\citep{Forveille09}. Excluding planets orbiting around pulsars, the first super-Earth detected was GJ 876d, a planet with $\\sim\\!\\!7.5 M\\E$ or\\nolinebreak[4]biting an M star \\citep{Rivera05}. Thereafter, other super-Earth planets have been discovered \\citep{Beaulieu06,Lovis06,Udry07,Ribas08,Mayor09,Forveille09,Howard09,Bouchy09}. In total, 12 planets with masses $\\lesssim 10 M\\E$ are reported today, and more detections are expected for the near future \\citep{Howard09,Mayor09}\\footnote{Up-to-date numbers can be found at the Extrasolar Planets Encyclopaedia: {\\tt http://www.exoplanet.eu}}. The least massive planet has a mass of 4.2 $M\\E$.  For M stars the habitable zone \\citep[HZ, defined by the range of orbital distances over which liquid water is possible on the planetary surface, see e.g.][]{Kasting93} is much closer to the star (depending on stellar mass, but typically $\\le$0.3 AU). Such close-in distances pose unique problems and constraints to habitability. For example, such planets are exposed to strong stellar winds in the early phases of the host star's evolution. Also, one expects a large number of coronal mass ejections (CMEs) on active M stars \\citep{Houdebine90,Scalo07}. The related interaction of the dense plasma flux (either from the stellar wind or a CME) with the atmosphere/magnetosphere environment of the exposed planets during the active stage of the stellar evolution could be strong enough to erode the atmosphere or the planets' water inventory via non-thermal atmospheric loss processes \\citep{Khodachenko07AB, Lammer07AB}. Weak magnetic shielding of tidally locked planets will also lead to an increased influx of galactic cosmic rays and stellar energetic particles \\citep{Griessmeier05a,Griessmeier09}, which can have important consequences for biological systems, both directly and indirectly via the modification of the planetary atmosphere \\citep[e.g.~destruction of atmospheric ozone, see][]{Grenfell07AB,Grenfell09}.  This paper is organised as follows: The evolution of the conditions in the stellar vicinity over the timescale of the stellar evolution is presented in Section \\ref{sec:stellarvar}. Two cases are considered: stellar CMEs (Section \\ref{sec:stellarvar:CME}), and the average stellar wind (Section \\ref{sec:stellarvar:sw}). Planetary magnetic shielding against the influence of the star is discussed in Section \\ref{sec:magneticshield}. Consquences are discussed in Section \\ref{sec:consequences}, with special focus on the potential evaporation of the planetary atmosphere (Section \\ref{sec:consequences:atm}) and the flux of cosmic rays to the planetary atmosphere (Section \\ref{sec:consequences:cr}). Section \\ref{sec:conclusions} closes with a few concluding remarks. ",
        "conclusions": "\\label{sec:conclusions}  Because of tidal locking, extrasolar planets with orbits in the habitable zone of K/M are likely to have much smaller magnetic moments than more distant planets. Combined with the strong ram pressure of dense and fast stellar winds and the high CME activity expected around young stars, this leads to the conclusion that such planets will have small magnetospheres which offer only limited magnetic protection.  The weak magnetic protection may lead to strong atmospheric erosion during the early stages of stellar evolution, which could be strong enough to erode the atmosphere or the planets' water inventory via non-thermal atmospheric loss processes. Another consequence of the weak magnetic protection is a high flux of galactic cosmic rays and stellar energetic particles to the planetary atmosphere. This has implications for potential habitability, but also for the atmospheric chemistry and composition (e.g.~destruction of atmospheric ozone). These effects have also to be considered for missions studying biosignatures in the observed spectra of Earth-like exoplanets.  Overall, the magnetic protection of exoplanets in the habitable zone of K/M stars can be weak, so that these planets may only be weakly shielded against stellar activity.  {\\footnotesize"
    },
    "1005/1005.0770_arXiv.txt": {
        "abstract": "We model the detectability of exoplanets around stars in the Beta Pic Moving Group (BPMG) using the Gemini Planet Imager (GPI), a coronagraphic instrument designed to detect companions by imaging.  Members of the BPMG are considered promising targets for exoplanet searches because of their youth ($\\sim 12$ MY) and proximity (median distance $\\sim 35$ pc). We wrote a modeling procedure to generate hypothetical companions of given mass, age, eccentricity, and semi-major axis, and place them around BPMG members that fall within the V-band range of the GPI.  We count as possible detections companions lying within the GPI's field of view and H-band fluxes that have a host-companion flux ratio placing them within its sensitivity. The fraction of companions that could be detected depends on their brightness at 12 Myr, and hence formation  mechanism, and on their distribution of semi-major axes.  We used brightness models for formation by disk instability and core-accretion.  We considered the two extreme cases of the semi-major axis distribution - the log-normal distribution of the nearby F and G type stars and a power-law distribution indicated by the exoplanets detected by the radial velocity technique.  We find that the GPI could detect exoplanets of all the F and G spectral type stars in the BPMG sample with a probability that depends on the brightness model and semi-major axis distribution.  At spectral type K to M1, exoplanet detectability depends on brightness and hence distance of the host star. GPI will be able to detect the companions of M stars later than M1 only if they are closer than 10 pc. Of the four A stars in BPMG sample, only one has V-band brightness in the range of GPI; the others are too bright. ",
        "introduction": "The primary goal of the Gemini Planet Imager (GPI), a coronagraphic instrument under construction for the Gemini Observatory (Graham et al. 2007; see also {\\it Future Instrumentation} at www.gemini.edu), is the detection of exoplanets directly by imaging.  Young stars will be important targets for GPI science because their exoplanets  are expected to dim as they age and cool. Over the past 20 years astronomers have identified groups of relatively nearby young stars distinguished by their space motion through the galaxy, hence the name \u201cNearby Young Moving Groups\u201d (NYMGs).  The TWA group at median distance 56 pc is the youngest, $\\sim 8$ MY, and the $\\beta$ Pic and AB Dor groups are the nearest with median distances 35 and 30 pc, respectively (Zuckerman and Song, 2004; Torres et al. 2006; Fern\\'andez et al. 2008). Their ages are about 12 MY for $\\beta$ Pic and 30-100 MY for AB Dor (Fern\\'andez et al. 2008). By analyzing the kinematics of the NYMGs, Fern\\'andez et al. (2008) showed that they share a common origin at the edges of the giant molecular cloud that formed the Sco-Cen star forming region. On this scenario, the supernova that heated the Local Bubble also triggered the formation of the stars in the NYMGs.  Owing to their youth and proximity stars in the NYMGs are promising targets for exoplanet searches by high resolution and high contrast imaging. With detectable contrast ratios of $10^7$, the GPI is designed to image brown dwarfs and exoplanets that may surround a star. Astronomers can then spectrally analyze the composition of their atmospheres, which has further implications in astrobiology and stellar and planetary evolution (eg. Oppenheimer and Hinkley, 2009).  L\\'epine and Simon (2009) described a pilot program to identify new low-mass probable members of the $\\beta$ Pic moving group (BPMG). Schlieder et al. (2009) are continuing the program and expanding it to include the AB Dor moving group.  Our goals in the work described here were to evaluate the detectability  of exoplanets in the BPMG and hence to provide guidance for observing programs of the BPMG, and other NYMGs, as new members are discovered. ",
        "conclusions": "We have found that for stars in the BPMG with $5<V<10$ mag,  1) GPI will be able to detect exoplanets with mass greater than $1 M_{J}$ for both hot start or core-accretion scenarios.  2) The V-band brightness for adequate AO correction requirement restricts the target sample to stars earlier than $\\sim$ K7 spectral type.  3) Given the GPI's high sensitivity, the two most important factors governing the exoplanet detection are the instrument's field of view and the period distribution of the exoplanets.  The number of exoplanets detected will distinguish between a distribution similar to that of the nearby F and G stars (i.e. DM91) and the power-law distribution derived by Cumming et al. (2008).   4) Stars $\\sim 10$ to 100 MY old are prime targets for GPI observation because of the favorable star-planet contrast.  For older stars such as a 1-Byr F star at the 35 pc  typical distance of the BPMG, the  GPI will be able to detect all companions with $M> \\sim 80 M_{J}$."
    },
    "1005/1005.5013_arXiv.txt": {
        "abstract": "{The iron lines at 630.15 and 630.25~nm are often used to determine the physical conditions of the solar photosphere. A common approach is to invert them simultaneously under the Milne-Eddington approximation. The same thermodynamic parameters are employed for the two lines, except for their opacities, which are assumed to have a constant ratio. } {We aim at investigating the validity of this assumption, since the two lines are not exactly the same.} {We use magnetohydrodynamic simulations of the quiet Sun to examine the behavior of the ME thermodynamic parameters and their influence on the retrieval of vector magnetic fields and flow velocities.}  {Our analysis shows that the two lines can be coupled and inverted simultaneously using the same thermodynamic parameters and a constant opacity ratio. The inversion of two lines is significantly more accurate than single-line inversions because of the larger number of observables. } {} ",
        "introduction": "\\label{sec:intro}  Both solar physics and radiative transfer owe many things to the Milne-Eddington (ME) approximation. Being an incredibly simplistic description of spectral line formation, it provides useful hints to understand the behavior of spectral lines as formed in solar and stellar atmospheres. It also offers a key diagnostics to infer the physical conditions of the plasma. This is particularly true when magnetic fields are present. The first solution of the radiative transfer equation in the presence of a magnetic field was derived adopting the ME assumption that all physical quantities relevant to line formation are constant with depth \\citep{1956PASJ....8..108U,1962aIzvKrAO...27.148R,1962aIzvKrAO...28.259R}. Under these conditions, the solution is analytic and, therefore, by simply varying the model parameters, one gets a handle on the behavior of Stokes line profiles. Similarly, the analytic character of the ME solution allows perturbative analyses like those performed by \\cite{1983SoPh...87..221L}, who studied the influence of velocity gradients, or \\cite{orozco2007}, who calculated the sensitivities of ME Stokes profiles to the various model parameters. Inversion codes of the radiative transfer equation are diagnostic tools that become simpler under the ME hypothesis. Since the pioneering work by \\cite{1972lfpm.conf..227H} and \\cite{1977SoPh...55...47A}, and after improvements by \\cite{1982SoPh...78..355L} and \\cite{1984SoPh...93..269L}, a number of ME inversion codes have been developed. These include the HAO code \\citep{lites2,1990ApJ...348..747L}, MELANIE \\citep{socas2001}, HELIX \\citep{2004A&A...414.1109L}, MILOS \\citep{orozco2007}, and VFISV \\citep{2010SoPh..tmp...35B}.  Strictly speaking, the ME model is applicable to just one spectral line. The reason is that the so-called thermodynamic parameters of the model are meant to characterize the behavior of the specific line under consideration. The line-to-continuum opacity ratio, $\\eta_{0}$, the Doppler width of the line, $\\Delta \\lambda_{\\rm D}$, and the damping parameter, $a$, govern the shape of the Stokes profiles (i.e., they are the parameters of the Voigt and Faraday-Voigt functions). In turn, the source-fuction terms, $S_{0}$ and $S_{1}$, control the continuum level, the line depression, and the Stokes amplitudes. However, many investigations are based upon the simultaneous ME inversion of spectral line pairs, like the well-known \\ion{Fe}{i} doublet at 630~nm (e.g., Lites et al. 1993) or the \\ion{Mg}{i}~{\\em b} lines at 517.2 and 518.3~nm \\citep{1988ApJ...330..493L}. These inversions are reported to use no extra free parameters as compared to single-line inversions, based on the similarities between the two lines belonging to the same multiplet: the opacities for each line are specified in the ratio of their respective oscillator strengths, while the Doppler widths and the damping parameters are assumed to be identical for the two lines. $S_{0}$ and $S_{1}$ are also the same for both lines since they lie very close in wavelength.  This strategy has been widely used for more than 20 years to analyze the \\ion{Fe}{i} 630 nm measurements taken with instruments such as the Advanced Stokes Polarimeter \\citep{1992SPIE.1746...22E} or the spectropolarimeter aboard the {\\em Hinode} satellite \\citep{2007ASPC..369...55L, 2007SoPh..243....3K, 2008SoPh..249..167T}. The simultaneous inversion has been shown to provide better results than single-line inversions \\citep[e.g.,][]{1994SoPh..155....1L}. The better performance is easy to understand: if the observables reproduced by the inversion are doubled, the results can be expected to be more accurate (at least by a factor $\\sqrt{2}$). However, the two lines do not have exactly the same atomic parameters (see Table\\ \\ref{cap5:tablelines}). This makes them to be formed at slightly different heights, where different physical conditions may exit \\citep[e.g.][]{2006A&A...456.1159M}.  In view of these differences and of the large excursions of $\\eta_{0}$, $a$, and $\\Delta \\lambda_\\mathrm{D}$ in the real solar photosphere, both horizontally and vertically, one may wonder whether the simultaneous ME inversion of the two lines is valid and what the limitations are. General inversion codes not relying on the ME approximation perform exact line transfer calculations, so they are able to invert the two lines without inconsistencies. The purpose of the present Research Note is to check the ME case using the excellent test bench offered by modern magnetohydrodynamic (MHD) simulations.  \\begin{table} \\caption{\\label{cap5:tablelines}Atomic data for the \\ion{Fe}{i} 630~nm lines. } \\centering \\begin{tabular}{@{}c c c c c c c@{}} \\hline $\\lambda_{0}$ (nm) & $\\chi$ (eV) & $\\log gf$ &  {\\scshape Transition }& $\\alpha/a^2_0$ & $\\sigma$ & $g_\\mathrm{eff}$\\\\ \\hline \\hline 630.1501 & 3.654 &  $-0.75$  & $5P_2-5D_2$ & 0.243 & 840.5 & 1.67 \\\\ 630.2494 & 3.686 &  $-1.236$ & $5P_1-5D_0$ & 0.240 & 856.8 & 2.5 \\\\ \\hline \\end{tabular} \\tablefoot{ Shown are the central wavelength of the transition, $\\lambda_{0}$, the excitation potential of the lower atomic level, $\\chi$, the multiplicity of the lower level times the oscillator strength, $\\log gf$, the collisional broadening parameters, $\\alpha$ and $\\sigma$ (in units of Bohr's radius $a_0$), and the effective Land\\'e factor of the line, $g_\\mathrm{eff}$. The $\\log gf$ values have been derived from a fit to the solar spectrum using a two-component model of the quiet Sun \\citep{borrero2002}. } \\end{table} ",
        "conclusions": ""
    },
    "1005/1005.3918_arXiv.txt": {
        "abstract": "The work presented in this paper aims at restricting the input parameter values of the semi-analytical model used in \\galics and \\momaf, so as to derive which parameters influence the most the results, \\eg star formation, feedback and halo recycling efficiencies, \\etc Our approach is to proceed empirically: we run lots of simulations and derive the correct ranges of values. The computation time needed is so large, that we need to run on a grid of computers. Hence, we model \\galics and \\momaf execution time and output files size, and run the simulation using a grid middleware: \\diet. All the complexity of accessing resources, scheduling simulations and managing data is harnessed by \\diet and hidden behind a web portal accessible to the users. ",
        "introduction": "\\label{sec:introduction}   Cosmological simulations in this context are used to simulate the evolution of dark matter through cosmic time in various universes. A classical simulation begins with an N-body computation using for example \\ramses~\\cite{2002AA_385_337T} or \\gadget~\\cite{2005MNRAS_364_1105S}. The output of the simulation is then post-processed using semi-analytical models such as \\galics~\\cite{2003MNRAS_343_75H} (\\textbf{GAL}axies \\textbf{I}n \\textbf{C}osmological \\textbf{S}imulations), then mock catalogs of observed galaxies are produced using \\momaf~\\cite{2005MNRAS_360_159B} (\\textbf{Mo}ck \\textbf{Ma}p \\textbf{F}acility). Those models use as input a set of parameters, which influence the results, as for example: the galaxy luminosity function, the number of galaxies per observed cones, and more globally the history of star formation rate in the galaxy evolution process. Those parameters have a large range of values, the experiment here aims at reducing for each parameter the intervals of values to a subset of ranges within which the output simulations would be coherent with observed galaxies distributions.  In order to stress how realistic the post-processing results are, a first important step is to identify the \\galics and \\momaf parameters which have the largest impact on the astrophysical results, such as star formation efficiency, feedback efficiency, halo recycling efficiency. As the parameters have a large range of possible values, exploring ``all'' combinations is really time consuming and requires lots of computing power. To harness the difficulty of running these analysis, one need firstly to have access to many computing resources, and secondly an efficient way to access those latter. Hence, we propose a client/server implementation for running post-processings on a distributed platform composed of heterogeneous machines: a \\emph{Grid}. A transparent access to these machines is provided by a \\emph{grid middleware}: \\diet (\\textbf{D}istribued \\textbf{I}nteractive \\textbf{E}ngineering \\textbf{T}oolbox). \\diet handles in one common and effective way the deployment of the computations on a grid of heterogeneous and distributed computers, the management of the different components of the post-processing, and also provides monitoring, communications and computation scheduling, and data and workflows management.  Running efficiently the post-processings on a set of distributed and heterogeneous machines requires estimations on both the execution time and the amount of data to be transfered for each applications. Hence, we benchmarked \\galics and \\momaf and derived their execution time and output file size. Those models are used within \\diet to select which computer should run the post-processing.   We first present in Section~\\ref{sec:cosmo} the cosmological simulation post-processing workflow, then we give the execution time and output file size models in Section~\\ref{sec:model}. In Section~\\ref{sec:ArchDIET} we give an overview of the \\diet middleware: its architecture, and the various features used within this project. Finally, and before concluding the paper, we present in Section~\\ref{sec:implem} the client/server implementation that allows the transparent execution of the post-processing workflow. ",
        "conclusions": "\\label{sec:conclusion}   In this paper, we presented models for \\galaxymaker's and \\momaf's execution time, and for \\galaxymaker's output files's size.  Execution time models give an estimation that is quite close to the real execution time for input files having a reasonable size. For small input files, the model returns an overestimate of the execution time, but this is easily explained by cache mechanisms at the hardware level, and can be taken into account in the model.  We also modeled the post-processing workflow, and provided efficient means of executed it in a parameter sweep manner on a grid of computers. Our solution relies on the \\diet middleware that provides transparent access to resources, and data and workflows management. All the complexity of running the post-processing on a grid, and the creation of the workflows are hidden by a web interface, that provides an easy and user-friendly way of submitting such cosmological simulations post-processing on many heterogeneous and distributed resources.  The next step is of course to analyze the results produced by the post-processing, and derive which parameters are the most influent on the results. Finally, we aim at providing ranges of values for the different parameters, that would provide correct results, \\ie comparable to observational data.      \\begin{theacknowledgments}  This work was developed with financial support from the ANR (Agence Nationale de la Recherche) through the LEGO project referenced ANR-05-CIGC-11.  \\end{theacknowledgments}"
    },
    "1005/1005.1406_arXiv.txt": {
        "abstract": "{ We perform spectral simulations of dynamo for magnetic Prandtl number of one with Taylor-Green forcing. We observe dynamo transition through a supercritical pitchfork bifurcation.   Beyond the transition, the numerical simulations reveal complex dynamo states with windows of constant, periodic, quasiperiodic, and chaotic magnetic field configurations. For some forcing amplitudes, multiple attractors were obtained for different initial conditions.  We show that one of the chaotic windows follows the period-doubling route to chaos.}   \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.4479_arXiv.txt": {
        "abstract": "For finite chemical potential effective models of QCD predict a first order phase transition. In favour for the search of such a phase transition in nature, we construct an equation of state for strange quark matter based on the MIT bag model. We apply this equation of state to highly asymmetric core collapse supernova matter with finite temperatures and large baryon densities. The phase transition is constructed using the general Gibbs conditions, which results in an extended coexistence region between the pure hadronic and pure quark phases in the phase diagram, i.e. the mixed phase. The supernovae are simulated via general relativistic radiation hydrodynamics based on three flavor Boltzmann neutrino transport in spherical symmetry. During the dynamical evolution temperatures above \\(10\\) MeV, baryon densities above nuclear saturation density and a proton-to-baryon ratio below \\(0.2\\) are obtained. At these conditions the phase transition is triggered which leads to a significant softening of the EoS for matter in the mixed phase. As a direct consequence of the stiffening of the EoS again for matter in the pure quark phase, a shock wave forms at the boundary between the mixed and the pure hadronic phases. This shock is accelerated and propagates outward which releases a burst of neutrinos dominated by electron anti-neutrinos due to the lifted degeneracy of the shock-heated hadronic material. We discuss the radiation-hydrodynamic evolution of the phase transition at the example of several low and intermediate mass Fe-core progenitor stars and illustrate the expected neutrino signal from the phase transition. ",
        "introduction": "The investigation of the QCD phasediagram via heavy-ion experiments at BNL's RHIC, CERN's LHC and the FAIR facility at GSI poised to explore the properties of QCD matter under extreme conditions. Three of the most important aspects are the behaviour and the position of the critical points in the phase diagram, the phase transition from hadronic matter to quark matter at finite chemical potentials and the properties of the predicted quark phases. In search for these aspects, observations of astronomical objects and astrophysical processes that are assumed to contain quark matter could be helpful. In such favour, Cold and isolated or binary neutrons stars (NS) have long been served as powerful objects to probe the equation of state (EoS) for hadronic as well as for quark matter. In the latter case the NS is called a hybrid star if in addition to the quark core a hadronic envelope is present. The astrophysical processes that leave an isolated NS as the final remnant are core collapse supernovae of low and intermediate mass Fe-core progenitor stars. Naturally the question rises at which stage during the dynamical evolution from a collapsing Fe-core to an isolated NS the thermodynamic conditions are such that the appearance of quark matter is favoured? Even further in the case of a phase transition from hadronic matter to quark matter, which hydrodynamic evolution can be expected and is there a relation to observations? The two intrinsically different scenarios are either during the early post bounce phase when the temperatures are moderately high or during the cooling of the remnant, where deleptonisation causes a temperature decrease and a density increase. The present article discusses the first case due to the critical conditions given by the quark matter EoS, where a relation to the explosion mechanism is explored. Therefore, a detailed study of core collapse supernovae including radiation-hydrodynamics with spectral neutrino transport and a sophisticated EoS for hot and dense asymmetric matter is required to simulate the matter conditions accurately.  The first study of the QCD phase transition in core collapse supernovae was published by \\citet{TakaharaSato:1988}, suggesting a relation of the multi-peak neutrino signal from supernova 1987A (see \\citet{Hirata:etal:1988}) to the appearance of quark matter. Using general relativistic hydrodynamics, they modelled the phase transition via a polytropic EoS. Due to the absence of neutrino transport they could neither confirm nor exclude the suggested prediction of a neutrino signal from the phase transition. Additional microphysics was included in the study by \\citet{Gentile:etal:1993}, where general relativistic hydrodynamics is coupled to a description of deleptonisation during the collapse phase of a progenitor star. Applying the MIT-bag model for the description of the quark phase, they find the formation of a (second) shock wave as a direct consequence of an extended co-existence region between the hadronic phase and the quark phase with a significantly smaller adiabatic index. The second shock wave follows and merges with the first shock from the Fe-core bounce after a few milliseconds. However, due to the lack of neutrino transport in the post bounce phase, they were also not able to confirm the predictions made for a possible neutrino signal from the phase transition. The recent investigation by \\citet{Nakazato:etal:2008} is based on general relativistic radiation hydrodynamics with spectral three flavour Boltzmann neutrino transport. They investigate very massive progenitors of \\(\\sim 100\\) M\\(_\\odot\\) which collapse to a black hole. Applying the MIT-bag model for quark matter, the time after bounce for black hole formation is shortened and corresponds to the appearance of quark matter, where the central mass exceeds the maximum stable mass given by the quark EoS.  We follow a similar approach and apply the MIT-bag model for the description of quark matter in general relativistic radiation hydrodynamics simulations, based on spectral three flavor Boltzmann neutrino transport. Our simulations are launched from low and intermediate mass progenitors, where no explosions could be obtained in spherical symmetry. We investigate the dynamical effects and discuss the possibility of an observable in the neutrino signal related to the QCD phase transition.  The manuscript is organised as follows. In \\S 2 we will present the standard core collapse supernova scenario and in \\S 3 we will lay down our neutrino radiation hydrodynamics model including both the hadron and quark EoSs. In \\S4 we will discuss the appearance of quark matter during the early post bounce phase of core collapse supernovae of intermediate mass Fe-core progenitors and summarise the results in \\S 5. ",
        "conclusions": "\\begin{table}[htpb] \\begin{center} \\caption{Selected properties of the PNSs for the different progenitor models under investigation. } \\begin{tabular}[b]{ c c c c c c c c } \\hline \\hline Prog. &  bag constant & \\(t_{pb}\\) & \\(M_\\text{PNS}\\) & \\(E_\\text{expl}\\) & \\(\\rho_\\text{c}\\) & \\(T_\\text{c}\\) & \\(Y_e\\) \\\\ \\([\\)M\\(_\\odot]\\)& \\([\\)MeV\\(]\\) & \\([\\)ms\\(]\\) & \\([\\)M\\(_\\odot]\\) & \\([10^{51}\\) erg\\(]\\) & \\([10^{14}\\) g/cm\\(^3]\\) & \\([\\)MeV\\(]\\) & \\\\ \\hline 10 & 162 & 248 & 1.434 & 0.361 & 6.6069 & 13.14 & 0.2343 \\\\ 10 & 165 & 429 & 1.482 & 1.080 & 5.8884 & 15.33 & 0.2488 \\\\ 13 & 162 & 241 & 1.467 & 0.146 & 6.4565 & 13.32 & 0.2335 \\\\ 13 & 165 & 431 & 1.498 & 0.323 & 5.6234 & 15.55 & 0.2462 \\\\ 15 & 162 & 209 & 1.608 & 0.420 & 6.7608 & 14.10 & 0.2262 \\\\ 15 & 165 & 330\\(^1\\) & 1.700 & unknown\\(^2\\) & 5.4954 & 15.33 & 0.2479 \\\\ \\hline \\end{tabular} \\end{center} \\small \\(^1\\) time of black hole formation\\\\ \\(^2\\) black hole formation before positive explosion energies are achieved \\label{table:quark-runs} \\end{table} The results of this first investigation of the QCD phase transition in core collapse supernova simulations of low and intermediate mass Fe-core progenitor stars is summarised in Table~\\ref{table:quark-runs}. The simulations are launched from progenitor stars of \\(10\\), \\(13\\) and \\(15\\) M\\(_\\odot\\) from \\citet{Woosley:etal:2002}. The post bounce times \\(t_{pb}\\) correspond to the appearance of the second neutrino burst in the spectra and the PNS masses \\(M_\\text{PNS}\\) are taken at the electron-neutrinospheres at some late post bounce times when the simulations are stopped. The central thermodynamic conditions, density \\(\\rho_\\text{c}\\), temperature \\(T_\\text{c}\\) and electron fraction \\(Y_e\\), correspond to the initial PNS collapse. A special model is the \\(15\\) M\\(_\\odot\\) progenitor using \\(B^{1/4}=165\\) MeV, where the maximum mass is reached shortly after the QCD phase transition. Hence, the PNS collapses to a black hole before the already formed second shock is accelerated to positive velocities. Due to our co-moving coordinate choice, no stable solution for the equations of energy and momentum conservation can be obtained and \\(t_{pb}\\) determines the time of black hole formation after bounce.  The significant softening of the EoS for matter which is in the mixed phase causes the PNS to collapse. As a direct consequence of the softening of the EoS for matter which reaches the pure quark phase, a second shock wave forms. The second shock accelerates and finally merges with the first shock and launches an explosion. The explosion energies \\(E_\\text{expl}\\) are give in Table~\\ref{table:quark-runs}. This mechanism was explored and found to even produce explosions in spherical symmetry using general relativistic radiation hydrodynamics based on spectral three-flavour Boltzmann neutrino transport, where otherwise no explosion could be obtained. The explosion energies in Table~\\ref{table:quark-runs} are evaluated at some late post bounce times when the simulations are stopped and might have not yet converged to the final value of the explosion energy. The simulations will have to be carried out longer. However, moderate explosion energies of \\(\\simeq 10^{51}\\) erg could be found for the \\(10\\) M\\(_\\odot\\) model using \\(B^{1/4}=165\\) MeV (i.e. the later phase transition), otherwise lower explosion energies are obtained. The explosion energies might be shifted to larger values when multi-dimensional effects such as convection and rotation are taken into account.  The second shock forms in the high density regime where neutrinos are fully trapped. Hence, no direct observational identification of the QCD phase transition can be found in the emitted neutrino spectra. It would be different if analysing the emitted gravitational waves directly from the phase transition (\\citet{Abdikamolov:etal:2009}). Unfortunately, gravitational waves have proven difficult to detect. Nevertheless, an indirect observation can be found in the emitted neutrino spectra. A second neutrino burst is released when the second shock, formed during the PNS collapse, propagates over the neutrinospheres. This second burst is, due to the lifted degeneracy of the shock heated hadronic material, dominated by electron-antineutrinos. The burst is accompanied by a significant increase of the mean neutrino energies which might become resolvable by neutrino detector facilities such as Super-Kamiokande for a future Galactic core collapse supernova explosion if quark matter appears.  One of the most important observations from supernova explosions is the composition of the ejecta, which is determined by explosive nucleosynthesis investigations and are model dependent (see for example Fr\\\"ohlich et al. (2006a-c)). Of special interest is the production of heavy elements for which rapid neutron captures (\\(r\\)-process) in supernova explosion models have long been investigated (see for example \\citet{WoosleyBaron:1992}, \\citet{Woosley:etal:1994}, \\citet{Witti:etal:1994}, \\citet{Otsuki:etal:2000}, \\citet{Thompson:etal:2001}, Wanajo et~al.(2006a,b), \\citet{Arcones:etal:2007}, \\citet{PanovJanka:2009}). The abundances are calculated via post processing of mass trajectories from explosion models and compared with the well known solar abundances. Unfortunately, the very recent neutrino driven explosion models fail to provide conditions favourable for the \\(r\\)-process, which are high entropies per baryon, a fast expansion timescales and generally neutron rich material. Especially the latter aspect is a subject of research. Although the explosion models achieved via the QCD phase transition of the present article have to be analysed consistently with respect to explosive nucleosynthesis, a reasonable amount of ejected matter is found to be neutron-rich where \\(Y_e\\simeq0.35-0.45\\). In addition, moderate entropies per baryon and a fast expansion timescale are obtained. The conditions found may indeed be favourable for the \\(r\\)-process."
    },
    "1005/1005.1892_arXiv.txt": {
        "abstract": "A large number of spectroscopic studies have provided evidence of the presence of multiple populations in globular clusters by revealing patterns in the stellar chemical abundances.  This paper is aimed at studying the origin of these abundance patterns.  We explore a model in which second generation (SG) stars form out of a mix of pristine gas and ejecta of the first generation of asymptotic giant branch stars.  We first study the constraints imposed by the spectroscopic data of SG stars in globular clusters on the chemical properties of the asymptotic and super asymptotic giant branch ejecta.  With a simple one-zone chemical model, we then explore the formation of the SG population abundance patterns focussing our attention on the Na-O, Al-Mg anticorrelations and on the helium distribution function.  We carry out a survey of models and explore the dependence of the final SG chemical properties on the key parameters affecting the gas dynamics and the SG formation process. Finally, we use our chemical evolution framework to build specific models for NGC 2808 and M4, two Galactic globular clusters which show different patterns in the Na--O and Mg--Al anticorrelation and have different helium distributions.  We find that the amount of pristine gas involved in the formation of SG stars is a key parameter to fit the observed O--Na and Mg--Al patterns. The helium distribution function for these models is in general good agreement with the observed one. Our models, by shedding light on the role of different parameters and their interplay in determining the final SG chemical properties, illustrate the basic ingredients, constraints and problems encountered in this self-enrichment scenario which must be addressed by more sophisticated chemical and hydrodynamic simulations. ",
        "introduction": "\\label{sec:introduction} In the last decade, a large number of spectroscopic and photometric studies have provided strong indication of the presence of multiple stellar populations in globular clusters.  The spectroscopic evidence comes from the observed spreads and anticorrelations between the abundances of light elements, not shown by the halo stars of similar metallicities. Specifically, all the clusters studied have shown an anticorrelation between Na and O and an anticorrelation between Mg and Al has been found in several massive clusters \\citep[cf.][]{grasneca04}. The fact that these chemical anomalies have been found also in turnoff and subgiant stars \\citep{gratton2001, ramirezcohen2002} indicates that they must be already present in the gas from which these stars formed. The anticorrelation between CN and CH bands observed in giant \\citep[e.g.][]{norrisfd1984} and turnoff stars indicates a difference in the nitrogen abundance and is considered another fingerprint of the presence of multiple populations in globular clusters.  An additional important constraint on the possible sources of gas from which second generation stars formed is provided by the observed constancy of the total C+N+O. The only exception, so far, is NGC 1851 for which a difference of ~0.6 dex is found among four red giants \\citep{yong2009} and for which the analysis of the subgiant branch suggests that the total C+N+O increases by about a factor three between the stars in the brighter and those in the dimmer branch \\citep{ventura18512009}.  Photometric studies have also provided a number of important results indicating the presence of significant differences in the helium abundances of stars within individual clusters. Indeed, the presence of helium differences seems to be the only explanation for the multiple main sequences or spread in the main sequence (MS) observed in some clusters. These observations confirmed the prediction of the presence of a helium spread based on the complex morphology of the horizontal branches of some clusters made by \\citep{daca04,dc2008}. In extreme cases (the blue MS of \\ocen\\ and of NGC 2808) a helium content Y=0.38--0.40 is inferred \\citep{norris2004,dantona2005,pio07}.  A recent spectroscopic survey by \\citet{carretta2009a} has  found that all the 15 clusters studied show the spectroscopic evidence of the presence  of multiple populations and that in all cases second generation stars represent a significant fraction  (50-80 \\%) of a cluster population. This implies that matter forming more than half of each cluster stars has been processed through the hot CNO cycle and by other proton-capture reactions on light nuclei.  In spite of the light elements variation, most GCs are mono--metallic objects, as far as abundances of heavier elements are concerned \\citep[see][for a recent review on this subject]{grasneca04}. Their heavy (Z $>$13) element metallicity, usually represented by the ratio [Fe/H], is found to be extremely homogeneous from star to star in each cluster. \\cite{carretta2009c} find that the upper limit to the scatter of iron is less than 0.05~dex in the 19 clusters they examine. This is consistent with previous determinations, e.g. for NGC~6752 \\citep[$\\sigma$=0.02][]{yong2005}, and for the clusters examined by \\cite{kraftivans2003} (0.03 -- 0.10~dex). The width in color of the main sequence and/or of the red giant branch also agree with the spectroscopic determination \\citep{suntzeff1993}.  As the site of production of heavy elements ($\\alpha$--capture and Fe--group elements) are stars exploding as core-collapse or thermonuclear supernovae \\citep[e.g.][]{wheeler1989}, the star to star iron homogeneity means that supernova ejecta do not affect the chemistry of the gas processed through the hot CNO cycle. On the other hand, significant iron and s--process elements spread is well known to be present in \\ocen\\ \\citep[e.g., for a recent analysis,][and references therein]{johnson2009} and more recently it has been confirmed in other massive clusters like M22 \\citep[][and references therein]{Marino09, dacosta2009}, M54 \\citep[][and references therein]{sl1995, bellazzini2008}, and Terzan~5 \\citep{ferraro2009}. In smaller clusters like NGC~1851, the SG might be enriched in calcium, according to \\cite{Lee09}, but see the discussion by \\cite{carretta2010}.  In \\citet{der08} (hereafter Paper I), we presented a model for the formation and dynamical evolution of multiple populations in globular clusters (GCs). In particular, we explored a model in which second generation (hereafter SG) stars form out of the ejecta of the first generation (hereafter FG) stars. By means of hydrodynamic simulations we have shown that the ejecta of AGB (asymptotic giant branch) stars collect in a cooling flow into the cluster core, where they form a subsystem of SG stars initially strongly concentrated in the cluster innermost regions. By means of N-body simulations we have then explored the subsequent stellar dynamical evolution of the cluster focusing our attention on the early loss of FG stars, on the evolution of the ratio of the number of FG and SG stars and on the evolution of their relative spatial distribution and mixing.  Paper I also included a preliminary discussion on the chemical abundances of FG AGB stars and on the possibility of reproducing the helium excess and other chemical anomalies observed in SG stars. Several theoretical and observational papers have shown that dilution of the FG ejecta with gas having the pristine composition is needed to explain the general shape of the anticorrelations \\citep{prantzos2007, bekki07,dantonaventura2007,carretta2009a}. If the SG anomalies are inputed to processing by Hot Bottom Burning (HBB) at the bottom of the convective envelope of massive AGBs, these models produce a direct correlation between sodium and oxygen abundance in the processed matter. This direct correlation is unavoidable, as found in all the relevant AGB computations \\citep{dh2003,herwig04,karakas2007,ventura2008a} and as explained in more detail in Section 2. While the total sodium yield is dependent on many uncertain factors, it is certain that a small oxygen depletion (lower burning temperature) goes together with a larger sodium abundance and a large oxygen depletion (higher burning temperature) is accompanied by a smaller sodium abundance. Any attempt to reverse this direct correlation, in the limits provided by available cross sections, has been unfruitful \\citep[e.g.][]{ventura2008a}. An observational hint in favour of the above arguments is given by \\citet{carretta2009a} who find a direct correlation between the minimum O and the maximum Na abundances of the 15 clusters they studied; such a result could not be understood if O and Na of the AGB ejecta were anticorrelated.  From all the above, it is apparent that, in order to explain the GC chemical patterns, it is necessary to introduce a form of dilution of the AGB ejecta, either with pristine gas, or with gas not showing the peculiar abundance patterns of the hot--CNO processed matter.  The source of this pristine matter still requires a detailed understanding.  \\cite{der08} proposed that initial asymmetries in the gas distribution allow to vent out the Type II supernovae (SNe~II) ejecta along preferential directions, creating an ``hour-- glass\" cavity and leaving some pristine gas in a torus at the outskirts of the cluster; at the end of the SN II activity this torus eventually collapses back into the cluster. In this scenario, the torus is not affected by contamination from supernova ejecta, and the iron content of the diluting matter is still the pristine one, consistently with the absence of iron differences between the FG and SG summarized above.  While the evolution of the larger clusters is probably also affected by more prolonged star formation and supernovae type Ia contamination \\citep[e.g.][]{mar07}, the presence of slight metal increase in smaller clusters might be a possible indication that the ``pristine\" gas in the falling back torus has been partially contaminated by the SN~II ejecta.  This paper is aimed at significantly expanding the initial study of SG stars chemical abundances presented in Paper I and at exploring in detail the SG chemical anomalies emerging from a model in which SG stars form from a mix of AGB ejecta and gas with pristine chemical composition.  We will refer to the scenario presented in Paper I. Consequently, we need to associate to the temporal evolution of the FG clusters the chemical composition of the ejecta of super--AGB and then massive AGB stars (hereinafter simply referred to as AGB pollutors) that successively provide the matter used for building up second generation stars. The chemical composition of the H-rich matter in the envelopes of these stars is affected by the hot CNO cycle processing at the basis of the convective envelope of these giants (hot bottom burning, hereafter HBB) and possibly also by the third dredge up (3DU) occurring after each thermal pulse. Stellar winds and planetary nebula ejection deposit the matter of these envelopes into the interstellar medium (ISM), where it can contribute to the formation of the second generation.  Despite the significant effort devoted to the determination of the yields of the massive AGBs and the numerous attempts to quantify the uncertainties in these calculations \\citep{ventura2005a,ventura2005b,karakas2007}, there are still large differences in the results obtained by different groups and even by the same groups at different times and no general consensus has been reached.  Our aim is to focus on the chemical properties of the stellar envelopes and explore the constraints on these properties imposed by the spectroscopic data of SG stars in globular clusters. We have addressed this problem by means of a simple one-zone model in which the ISM is supplied by the FG AGB ejecta and by the possible accretion of pristine gas, and is depleted by the SG star formation (hereafter SF).  This simplified approach allows to easily control the evolution of the amount of the AGB ejecta and pristine gas involved in the SG formation process and to fully explore the chemical properties of the ISM and of the SG stars.  The structure of the paper is the following. We start in sect. \\ref{sec:yelds} with an overview of the chemical abundances of the massive AGB ejecta. We then illustrate our general assumptions in sect. \\ref{sec:assump} and the model details in sect. \\ref{sec:comp}. We then analyze the model in sect. \\ref{sec:moana}, and compare our results with the data of NGC 2808 and M4 in sect. \\ref{sec:ngc2808} and sect. \\ref{sec:m4}, respectively. Finally, we summarize our conclusions in sect. \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc} In this paper we have studied the origin of the chemical patterns which have been observed in many globular clusters and which are considered the spectroscopic fingerprints of the presence of multiple stellar populations.  Specifically, in our investigation we have focussed our attention on the O-Na and Mg-Al anticorrelations and the helium distribution function.  In our model we have assumed that AGB stars are the polluting source of gas with the anomalous abundances of light elements observed in globular cluster second-generation stars.  Our chemical framework is based on a one-zone model following the formation of SG stars from a mix of ejecta of AGB stars and pristine gas.  We have carried out a large number of simulations to explore the dependence of the SG chemical properties on the parameters characterizing the star formation process and the dynamics of the involved pristine gas. Finally, we have used our framework to model the observed chemical patterns in the Galactic clusters NGC 2808 and M4.  The main results of our study are the following:  \\begin{enumerate} \\item The current stellar models provide AGB ejecta in which both Na and O decrease with increasing stellar mass. In order to reproduce the observed O-Na anticorrelation, the gas from which SG stars form must be diluted with O-rich, Na-poor pristine gas (see e.g. Figs. \\ref{fig:refmod} and \\ref{fig:refchm}). \\item The amount of pristine gas involved in the SG formation process, the timescales driving the dynamics of such gas, and the star formation efficiency play a key role in determining the extension of the O-Na anticorrelation and the fraction of extreme Na-rich/O-poor stars. \\item The helium abundance distribution function is correlated with the distribution of stars in the O-Na plane. Extreme Na-rich/O-poor stars are also those with an extreme He enrichement. Our models predict that all the clusters with a very extended O-Na anticorrelation should also host a population of He-rich stars. \\item Our models show that the extension of the O-Na anticorrelation is closely correlated with that in the Mg-Al plane; Na-rich/O-poor stars have also high Al and low Mg abundances. \\item We have used our framework to build specific models for two prototypical Galactic globular clusters: NGC 2808, a massive cluster that hosts a SG population characterized by a very extended O-Na anticorrelation and includes a very He-rich population, and M4, a low-mass cluster with a significantly less extended O-Na anticorrelation that does not include extremely Na-rich O-poor stars and for which there is no photometric evidence of a He-rich population. Despite the significant differences in their chemical patterns, in both clusters a significant fraction of stars belong to the SG \\citep[50 per cent in NGC 2808,] [and 65 per cent in M4 \\cite{carretta2009a,mar08})]{carretta2009a}.   Our models successfully reproduce the differences in the O-Na anticorrelation observed in these two clusters, the distribution of stars in the Mg-Al plane in M4 and predict an extended Mg-Al anticorrelation for NGC 2808. Al and Mg abundances for NGC 2808 have been determined with UVES observations only for a small number of either Na-rich/O-poor or Na-poor/O-rich stars. As predicted by our model, these stars populate only the Al-rich/Mg-poor and the Al-poor/Mg-rich regions of the Mg-Al plane. According to our model, future determinations of Al and Mg abundances of targeting stars with intermediate Na and O should lead to intermediate Al and Mg abundances filling the extended Al-Mg anticorrelation.  \\item The helium distribution, including the extreme population formed directly from super--AGB ejecta  is well reproduced by our model for NGC~2808 (see Fig. \\ref{fig:2808}). In clusters like M4, in which a larger dilution with pristine matter is necessary to model the O-Na and Mg-Al patterns, a large helium dispersion is not required, but a small helium spread should generally be present (see Fig. \\ref{fig:m4}). \\end{enumerate}  Our investigation has shed light on the key chemical and hydrodynamical ingredients determining the formation of the chemical patterns observed in globular clusters. The results presented here are to be the starting point informing further study based on full hydrodynamical simulations. Several issues will require further investigation. Specifically, as for the stellar evolution models, we have shown that our conservative choice of an educated linear and monotonic extrapolation for the sodium and oxygen abundances in the mass range from 6.5 to 9 $\\msun$ reproduces the general trend of the O--Na anticorrelations of different clusters, when also the differences in neon abundances in the FG of different clusters are taken into account. In NGC~2808, however, we can not reproduce the (few) very large sodium values of some stars for which only upper limits on the oxygen abundance are available (see Fig. \\ref{fig:2808}). Another parameter is probably at work here, and only further investigation into the super--AGB phase will shed light on this issue.  The lack of any evidence of a significant metal enrichment in most clusters hosting multiple stellar populations implies that neither ejecta from FG nor SG supernovae  is involved in the chemical enrichment of matter from which SG stars form. While in our model SG formation starts after the end of the FG SN II epoch, it is to be further clarified whether the lack of metal enrichment from SG supernovae is due to SG forming with IMF truncated at $M<9$ $\\msun$ or whether this is a consequence of a more complex gas dynamics to be further explored with full hydrodynamical simulations.  Finally, additional full hydrodynamical simulations will be needed to clarify the source of pristine gas and the accretion mechanism. Different processes have been considered in the literature, such as accretion from a diffuse surrounding medium \\citep[e.g.][]{limu07,pflkro09} or interaction with molecular clouds \\citep{bekmac09}. Conditions for an effective gas collection within the cluster turn out to be rather specific. Instead, our model requires a mechanism general enough to work in all the clusters and able to take into account the cluster-to-cluster differences in the dynamics and the amount of pristine gas involved in the SG formation process."
    },
    "1005/1005.2497_arXiv.txt": {
        "abstract": "{The properties of the dust grains (e.g., temperature and mass) can be derived from fitting far-IR SEDs ($\\geq$100~\\micron).  Only with SPIRE on {{\\it Herschel}} has it been possible to get high spatial resolution at 200 to 500~\\micron\\ that is beyond the peak ($\\sim$160~\\micron) of dust emission in most galaxies.} {We investigate the differences in the fitted dust temperatures and masses determined using only $<$200~\\micron\\ data and then also including $>$200~\\micron\\ data (new SPIRE observations) to determine how important having $>$200~\\micron\\ data is for deriving these dust properties.} {We fit the 100 to 350~\\micron\\ observations of the Large Magellanic Cloud (LMC) point-by-point with a model that consists of a single temperature and fixed emissivity law.  The data used are existing observations at 100 and 160~\\micron\\ (from IRAS and {{\\it Spitzer}}) and new SPIRE observations of 1/4 of the LMC observed for the HERITAGE Key Project as part of the {{\\it Herschel}} Science Demonstration phase.} {The dust temperatures and masses computed using only 100 and 160~\\micron\\ data can differ by up to 10\\% and 36\\%, respectively, from those that also include the SPIRE 250 \\& 350~\\micron\\ data.  We find that an emissivity law proportional to $\\lambda^{-1.5}$ minimizes the 100--350~\\micron\\ fractional residuals.  We find that the emission at 500~\\micron\\ is $\\sim$10\\% higher than expected from extrapolating the fits made at shorter wavelengths.  We find the fractional 500~\\micron\\ excess is weakly anti-correlated with MIPS 24~\\micron\\ flux and the total gas surface density.  This argues against a flux calibration error as the origin of the 500~\\micron\\ excess.  Our results do not allow us to distinguish between a systematic variation in the wavelength dependent emissivity law or a population of very cold dust only detectable at $\\lambda \\geq 500~\\micron$ for the origin of the 500~\\micron\\ excess.} {} ",
        "introduction": "Among nearby galaxies, the Large Magellanic Cloud (LMC) and Small Magellanic Cloud (SMC) represent unique astrophysical laboratories for interstellar medium (ISM) studies.  Both Clouds are relatively nearby, the LMC at $\\sim$50 kpc \\citep{Schaefer08} and the SMC at $\\sim$60 kpc \\citep{Hilditch05}, and provide ISM measurements that are relatively unconfused by line-of-sight uncertainties when compared to the Milky Way.  The two Clouds span an interesting metallicity range with the LMC at $\\sim$1/2 Z$_{\\sun}$ \\citep{Russell92} being above the threshold of 1/3--1/4 Z$_{\\sun}$ where the properties of the ISM change as traced by the rapid reduction in the PAH dust mass fractions and possible dust-to-gas ratios \\citep{Draine07} and the SMC at $\\sim$1/5Z$_{\\sun}$ \\citep{Russell92} below this threshold. Finally, the dust in the LMC and SMC shows strong variations in its ultraviolet characteristics \\citep{Gordon03}.  The HERschel Inventory of The Agents of Galaxy Evolution (HERITAGE) in the Magellanic Clouds {\\it Herschel} Key Project will map both Clouds using the PACS/SPIRE Parallel observing mode providing observations at 100, 160, 250, 350, and 500~\\micron\\ \\citep{Meixner10}. The HERITAGE wavelength coverage (100--500~\\micron) and spatial resolution ($\\sim$10~pc at 500~\\micron) is well suited to measuring the spatial variations of dust temperatures and masses.  The infrared dust emission in most galaxies peaks between 100--200~\\micron\\ \\citep{Dale05} and observations $>$200~\\micron\\ are important for accurate dust temperature and masses \\citep{Willmer09}.  Ground-based submilimeter observations do provide the needed $>$200~\\micron\\ observations, but they have been seen to be in excess of that expected from extrapolating fits to the $<$200~\\micron\\ data for sub-solar metallicity galaxies \\citep{Galliano05}.  This excess could be due to very cold dust that only emits at submilimeter wavelengths or variations in the wavelength dependent dust emissivity law \\citep{Reach95, Paradis09}.  As part of the Science Demonstration Program (SDP), two HERITAGE AORs centered on the LMC were executed. These observations are used in this paper to explore the impact SPIRE observations have on the measurement of dust temperatures and masses including the behavior of any submilimeter excess. ",
        "conclusions": "We investigate the importance of $>$200~\\micron\\ data in determining dust temperatures and masses using new {\\it Herschel} SPIRE observations of the LMC (taken for the HERITAGE Key Project as part of the {\\it Herschel} Science Demonstration phase) combined with existing IRAS 100~\\micron\\ and {\\it Spitzer} MIPS 160~\\micron\\ images.  We fit the observations with a model consisting of dust emitting as a single temperature blackbody modified with an emissivity law proportional to $\\lambda^{-\\beta}$.  For fixed values of $\\beta$, fits using only the 100--160~\\micron\\ data give dust temperatures and masses that are on average up to 8\\% and 23\\% different from fits using the same $\\beta$ and the 100--350~\\micron\\ data.  The new SPIRE observations allowed us to determine that $\\beta = 1.5$ minimizes the residuals from 100 to 350~\\micron.  Using a $\\beta = 2.0$ for the 100--160~\\micron\\ and a $\\beta = 1.5$ for the 100--350~\\micron\\ fits results in an increase of 10\\% for the dust temperature and a decrease in the dust mass by 36\\%.  On average, there is a fractional excess at 500~\\micron\\ of $\\sim$10\\%.  The origin of the fractional excess is unlikely to be due to our fitting algorithm or a flux calibration error, but it could be due to either very cold dust that emits only $\\geq$500~\\micron\\ or a variation in the wavelength dependent change in the dust emissivity. Planned HERITAGE observations of the LMC and SMC will allow for a more detailed investigation of including $> 200~\\micron$ data (mainly the 500~\\micron\\ excess) due to better quality PACS and SPIRE images (optimized observations and cross-scans)."
    },
    "1005/1005.2342_arXiv.txt": {
        "abstract": "{Cyg OB2 \\#5 is a contact binary system with variable radio continuum emission. This emission has a low-flux state where it is dominated by thermal emission from the ionized stellar wind and a high-flux state where an additional non-thermal component appears. The variations are now known to have a period of 6.7$\\pm$0.2 yr. The non-thermal component has been attributed to different agents: an expanding envelope ejected periodically from the binary, emission from a wind-collision region, or a star with non-thermal emission in an eccentric orbit around the binary. The determination of the angular size of the non-thermal component is crucial to discriminate between these alternatives. We present the analysis of VLA archive observations made at 8.46 GHz in 1994 (low state) and 1996 (high state), that allow us to subtract the effect of the persistent thermal emission and to estimate an angular size of $\\leq 0\\rlap.{''}02$ for the non-thermal component. This compact size favors the explanation in terms of a star with non-thermal emission or of a wind-collision region. } ",
        "introduction": "Cyg OB2 \\#5 (V729 Cyg, BD +40 4220) is an eclipsing, contact binary system consisting of two O-type supergiants with a 6.6-day period (Hall 1974; Leung \\& Schneider 1978; Rauw et al. 1999; Linder et al. 2009). As several other luminous O-star systems in the Cyg OB2 association, this source was found to show evidence of variable radio emission (Persi et al. 1985, 1990; Bieging et al. 1989). The radio emission appears to have two states: a low-flux state of  $\\sim$2 mJy at 4.8 GHz where the spectral index is consistent with thermal emission from an ionized stellar wind, and a high-flux state of $\\sim$8 mJy at 4.8 GHz where the spectral index is flatter than in the low state. The variations were proposed to have a 7-year period (Miralles et al. 1994) and have been attributed to variable non-thermal emission from an expanding envelope arising periodically in the binary (Persi et al. 1990; Bieging et al. 1989; Miralles et al. 1994).  In addition, Abbott et al. (1981) and Miralles et al. (1994) reported on the existence of a radio ``companion'' 0.8 arcsec to the NE of the main radio source (which is associated with the eclipsing binary).  Observations by Contreras et al. (1997) revealed that this radio source has an elongated shape and lies in-between the short-period binary and a third star, which was first reported by Herbig (1967). Contreras et al. (1997) suggested that the proposed NE radio companion actually corresponds to the wind interaction zone between the binary system and the tertiary component. Recently, Kennedy et al. (2010) reanalysed all VLA observations of Cyg OB2 \\#5 and showed that the primary radio source, associated with the eclipsing binary, varies with a period of  $6.7 \\pm 0.2$ yr while the flux from the secondary NE source remains constant in time. These authors proposed that the variations in the main radio component can be represented by a simple model in which a fourth star (with constant non-thermal emission) moves around the eclipsing binary in an eccentric orbit and the varying radio emission results from the variable free-free opacity of the wind between the star and the observer (Kennedy et al. 2010). These authors estimate a major axis of 14 AU for the star with non-thermal emission. In this scenario, we have a quadruple star system: the contact binary, the star with non-thermal emission in a 6.7 yr orbit around the contact binary, and the NE component.  To advance our understanding of the nature of the time-variable non-thermal emission, a determination of the angular size of the source producing it is needed. Using VLA data taken at 8.46 GHz in 1991 October 3 (when Cyg OB2 \\#5 was in the high radio state), Miralles et al. (1994) estimated an angular size of $0\\rlap.{''}046 \\pm 0\\rlap.{''}006$ for the whole (non-thermal plus thermal components) emission. Using MERLIN data taken at 4.8 GHz in 1996 November 14 (when Cyg OB2 \\#5 was again in the high radio state), Kennedy et al. (2010) estimated an angular size of $\\sim0\\rlap.{''}077$ for the whole (non-thermal plus thermal components) emission associated with the primary radio component. Since the thermal contribution comes from the ionized wind that is known to be extended, these size estimates can be considered as upper limits to the size of the non-thermal emission. In this paper we present the analysis of archive 8.46 GHz continuum observations made with the Very Large Array (VLA) of the NRAO\\footnote{The National Radio Astronomy Observatory is operated by Associated Universities Inc. under cooperative agreement with the National Science Foundation.} in the A configuration toward Cyg OB2 \\#5 in two epochs. During 1994 April 9 (1994.27), the source was in the low radio state and only the thermal emission from the ionized stellar wind was present. These observations were used to estimate the characteristics of the persistent thermal component. During 1996 December 28 (1996.99), the source was in the high radio state, with both the thermal and non-thermal components present. Using the information obtained from the 1994 data, it is possible to subtract in the \\sl(u,v) \\rm plane the contribution from the thermal component and obtain a more stringent estimate of the angular size of the non-thermal emission. ",
        "conclusions": "We present the analysis of VLA archive data of the Cyg OB2 \\#5 system taken during the high (1996) and low (1994) radio states. This analysis allows the subtraction, in the \\sl (u,v) \\rm plane of the persistent thermal component and a better estimate of the angular dimensions of the non-thermal component. We obtain an upper limit, $\\theta('') \\leq 0\\rlap.{''}02$ for the angular size of the non-thermal emission. Cyg OB2 \\#5 will be again in high state during 2010 and 2011 and a very long baseline interferometry study is needed to determine the angular size and morphology of the compact non-thermal emission."
    },
    "1005/1005.1207_arXiv.txt": {
        "abstract": "{The question of how much gas cools in the cores of clusters of galaxies has been the focus of many, multiwavelength studies in the past 30 years. In this letter we present the first detections of the strongest atomic cooling lines, [C{\\sc ii}], [O{\\sc i}] and [N{\\sc ii}] in two strong cooling flow clusters, A1068 and A2597, using {\\it Herschel} \\rm PACS. These spectra indicate that the substantial mass of cold molecular gas ($>10^{9}$M$_\\odot$) known to be present in these systems is being irradiated by intense UV radiation, most probably from young stars. The line widths of these FIR lines indicate that they share dynamics similar but not identical to other ionised and molecular gas traced by optical, near-infrared and CO lines. The relative brightness of the FIR lines compared to CO and FIR luminosity is consistent with other star-forming galaxies indicating that the properties of the molecular gas clouds in cluster cores and the stars they form are not unusual. These results provide additional evidence for a reservoir of cold gas that is fed by the cooling of gas in the cores of the most compact clusters and provide important diagnostics of the temperature and density of the dense clouds this gas resides in.}  \\date{Received 30 March 2010/Accepted} ",
        "introduction": "The cooling process at the cores of galaxy clusters is highly complex: recent {\\it XMM-Newton} and {\\it Chandra} observations indicate that the cooling rates are reduced by an order of magnitude below the simple cooling flow models at temperatures below $\\sim2\\times10^7$K (Peterson \\& Fabian 2006). These X-ray observations, when linked with the detection of radio jet inflated bubbles in the cores of many of the strongest cooling flows (see McNamara \\& Nulsen 2007 for a review), suggest that the strong suppression of gas cooling is related to energy injection into the intracluster medium by the action of jets and related AGN activity.  The detection of substantial masses of molecular gas in the cores of the most rapidly cooling clusters through CO lines (Edge 2001, Salom\\'e \\& Combes 2003) and warm H$_2$ molecular lines in the NIR and MIR (Jaffe \\& Bremer 1997, Egami et al. 2006) indicates that not all cooling is suppressed and this cooled gas may provide the fuel for future AGN activity. These tracers of molecular gas appear to correlate with the strength of optical lines from ionised gas (Crawford et al. 1999, Edge 2001) and the dust continuum at MIR and sub-mm wavelengths (O'Dea et al. 2008). However, the excitation of these various emission lines and the relative importance of energy input from star formation, AGN, cosmic rays and/or the intracluster medium is poorly constrained (Ferland et al. 2009).  One as yet unexplored diagnostic of the properties of the cold gas are the atomic cooling lines found in the FIR, [C{\\sc ii}], [O{\\sc i}] and [N{\\sc ii}]. The unprecedented sensitivity of {\\it Herschel} \\rm (Pilbratt et al. 2010) to FIR line emission offers the opportunity to assess the ionisation and density of the colder gas for the first time with the [C{\\sc ii}] line and two principle [O{\\sc i}] lines. The authors were awarded 140~hours of time in an Open Time Key Project (PI Edge) to investigate the FIR line and continuum properties of a sample of 11 brightest cluster galaxies (BCGs) in well-studied cooling flow clusters selected on the basis of optical emission line and X-ray properties. The full goals of the project are to observe at least five atomic cooling lines for each object that cover a range in density and temperature behaviour and obtain a fully sampled FIR spectral energy distribution for systems where significant star formation is expected. In this paper we present the Photodetector Array Camera \\& Spectrometer (PACS, Poglitsch et al. 2010) spectroscopy for the two targets observed in the Science Demonstration Phase (SDP), Abell~1068 ($z=0.1386$) and Abell~2597 ($z=0.0821$). In a parallel paper (Edge et al. 2010), we present the FIR photometry for these clusters.  The two clusters observed have quite contrasting multiwavelength properties. A1068 is a strong MIR source (O'Dea et al 2008) with a bright CO detection (Edge 2001) but a weak radio source (McNamara et al 2004). A1068 lies just below the luminosity threshold of a ULIRG (10$^{12}$~L$_\\odot$) and exhibits some contribution from an AGN (Crawford et al. 1999, O'Dea et al. 2008). On the other hand, A2597 is a relatively weak MIR source (Donahue et al. 2007) with a weak CO detection (Salom\\'e, priv. comm.) but a powerful radio source (Sarazin et al. 1995). The implied FIR luminosity of A2597  of  8.8$\\times 10^{9}$~L$_\\odot$ is a factor of around 30 below that of A1068  (3.5$\\times 10^{11}$~L$_\\odot$) and,  in addition, the fractional contribution from an AGN in the MIR is also lower. ",
        "conclusions": "These initial results from {\\it Herschel} \\rm indicate that atomic cooling lines are present in the brightest cluster galaxies in cooling flow clusters. The intensity and velocity width of these lines is consistent with all the other observed tracers of cold gas in these systems implying they originate from the same population of clouds. The only apparent exception to this in our current observations is that the FIR lines appear to be systematically broader than the CO lines impling that the {\\it relative} intensity of these lines varies with position within the BCG. The results that will come from our Open Time Key Project for 11 BCGs will expand greatly on those presented here with more lines and a greater range of BCG properties. Beyond this, the potential for {\\it Herschel} \\rm to illuminate the properties of the cold gas that may fuel cold nuclear accretion in more distant clusters and local groups is vast."
    },
    "1005/1005.1031_arXiv.txt": {
        "abstract": "{During an X-ray survey of the Small Magellanic Cloud, carried out with the XMM-Newton satellite, we detected significant soft X-ray emission from the central star of the high-excitation planetary nebula SMP SMC 22. Its very soft spectrum is well fit with a non local thermodynamical equilibrium model atmosphere composed of H, He, C, N, and O, with abundances equal to those inferred from studies of its nebular lines. The derived effective temperature of 1.5$\\times10^5$ K is in good agreement with that found from the optical/UV data. The unabsorbed flux in the 0.1--0.5 keV range is $\\sim3\\times10^{-11}$ erg cm$^{-2}$ s$^{-1}$, corresponding to a luminosity of $\\sim1.2\\times10^{37}$ erg s$^{-1}$  at the distance of 60 kpc. We also searched for X-ray emission from a large number of SMC planetary nebulae, confirming the previous detection of SMP SMC 25 with a luminosity of (0.2--6)$\\times10^{35}$ erg s$^{-1}$ (0.1-1 keV). For the remaining objects that were not detected, we derived flux upper limits corresponding to luminosity  values from several tens to hundreds times smaller than that of SMP SMC 22. The exceptionally high X-ray luminosity of SMP SMC 22 is probably due to the high mass of its central star, quickly evolving toward the white dwarf's cooling branch, and to a small intrinsic absorption in the nebula itself.}  \\keywords {planetary nebulae: individual: SMP SMC 22, SMP SMC 25 - Magellanic Clouds - X-rays: general} ",
        "introduction": "Planetary nebulae (PNe) are a common stage in the evolution of low  and intermediate mass stars, leading to the formation of white dwarfs. They appear when the fast wind from the central star interacts with the matter of the denser wind that was previously ejected during the asymptotic giant branch (AGB) phase, and are characterized by  H$_{\\alpha}$ line emission. Since the advent of imaging X-ray telescopes a number of PNe have been detected in the soft X-ray range (see, e.g., \\citealt{chu03}), but due to their relatively low fluxes,   detailed studies with the \\xmm\\ and Chandra satellites have been carried out only for a few objects \\citep{gru06,kas07,mon09}. X-ray emission from PNe can originate from their central star or from the hot gas shocked in the interaction between the two stellar winds. Usually one of these two processes, characterized by different spectral and spatial signatures, is dominant, but there are also PNe in which both components have been detected. There is also the possibility that some X-ray emission observed from PNe is actually due to coronal emission from a companion star (see, e.g., \\citealt{sok02}).   SMP SMC 22 (hereinafter \\sm ) is a  high-excitation planetary nebula located in the Small Magellanic Cloud (SMC) \\citep{san78,all87} and  characterized by a very high X-ray luminosity \\citep{wan91}. Its large X-ray flux led to an  early detection of this source in the soft X-ray range with the Einstein Observatory \\citep{sew81} and to its inclusion in the class of super-soft X-ray sources (SSS), a rather heterogeneous group of luminous (10$^{36}$--10$^{38}$ erg s$^{-1}$) sources characterized by thermal-like emission corresponding to blackbody temperatures of $10^5$--10$^6$ K (see, e.g., \\citealt{kah06}). Most SSS are binary systems containing accreting white dwarfs, but a few of them have been identified with the nuclei of PNe. Here we report on recent  X-ray observations of \\sm\\ obtained with the \\xmm\\ satellite, as well as on a systematic search for X-ray emission from a large sample of PNe in the SMC. ",
        "conclusions": "The first X-ray observations of the SMC planetary nebula \\sm\\ obtained with a modern high-throughput satellite have allowed  us to study its X-ray emission with unprecedented statistics. No evidence for a binary nature, such as long or short term variability, as seen in other SSS was found. It is remarkable that, despite different spectral models can fit the data, a self-consistent picture in terms of temperature, mass and radius of the central star can be obtained with a NLTE model atmosphere with the same elemental abundances seen in the nebula. The inferred mass for the central star, of the order of 1 $\\msun$,   implies that \\sm\\ is the descendent  of a relatively massive progenitor (see Fig. \\ref{HR}). This may explain its exceptional luminosity,   as well as the apparent rarity of such objects that evolve very quickly toward the cooling white dwarfs sequence."
    },
    "1005/1005.4422_arXiv.txt": {
        "abstract": "With the aim of determining if Milky Way (MW) progenitors could be identified as high redshift Lyman Alpha Emitters (LAEs) we have derived the intrinsic properties of $z\\approx 5.7$ MW progenitors, which are then used to compute their observed Ly$\\alpha$ luminosity, $L_\\alpha$, and equivalent width, EW. MW progenitors visible as LAEs are selected according to the canonical observational criterion, $L_\\alpha>10^{42}$~erg~s$^{-1}$ and $EW>20$ \\AA. Progenitors of MW-like galaxies have $L_{\\alpha} =10^{39-43.25}$~erg~s$^{-1}$, making some of them visible as LAEs. In any single MW merger tree realization, typically only 1 (out of $\\approx 50$) progenitor meets the LAE selection criterion, but the probability to have {\\it at least} one LAE  is very high, $P=68\\%$. The identified LAE stars have ages, $t_* \\approx  150-400$~Myr at $z\\approx 5.7$ with the exception of five small progenitors with  $t_* < 5$~Myr and large $EW=60-130$~\\AA. LAE MW progenitors provide $> 10\\%$ of the halo very metal-poor stars [Fe/H]$<-2$, thus establishing a potentially fruitful link between high-$z$ galaxies and the Local Universe. ",
        "introduction": "\\label{intro} Lyman Alpha Emitters (LAEs) are galaxies identified by means of a very strong Ly$\\alpha$ line (1216 \\AA). Advances in instrument sensitivity and specific spectral signatures (strength, width and continuum break bluewards of the line) have enabled the confirmation of hundreds of LAEs in a wide redshift range, at $z\\approx 2.25$ (Nilsson et al. 2008), $z\\approx3$ (Cowie \\& Hu 1998; Steidel et al. 2000; Matsuda et al. 2005; Venemans et al. 2007; Ouchi et al. 2008), $z\\approx 4.5$ (Finkelstein et al. 2007), $z\\approx5.7$ (Malhotra et al. 2005; Shimasaku et al. 2006) and $z\\approx 6.6$ (Taniguchi et al. 2005; Kashikawa set al. 2006).  LAEs have by now been used extensively as probes of both the ionization state of the intergalactic medium (IGM) and probes of high redshift galaxy evolution (Santos 2004; Dayal, Ferrara \\& Gallerani 2008; Nagamine et al. 2008; Dayal et al. 2009; Dayal, Ferrara \\& Saro 2010; Dayal, Maselli \\& Ferrara 2010). However, in spite of the growing data sets, there has been no effort to establish a link between the properties of these early galaxies to observations of the local Universe, {\\it in primis} the Milky Way (MW). Our aim in this work is to investigate the possible connection between the Galactic building blocks and LAEs at a time when the Universe was $\\approx 1$ Gyr old. This will allow us to answer to questions as: are the progenitors of MW-like galaxies visible as LAEs at high redshifts? How can we discriminate amongst LAEs which are possible MW progenitors and those that are not? What are the physical properties of these Galactic building blocks?  To this end, we adopt a novel approach of coupling the semi-analytical code {\\tt GAMETE} (Salvadori, Schneider \\& Ferrara 2007; Salvadori, Ferrara \\& Schneider 2008, Salvadori \\& Ferrara 2009), which traces the hierarchical build-up of the Galaxy, successfully reproducing most of the observed MW and dwarf satellite properties at $z=0$, to a previously developed LAE model (Dayal, Ferrara \\& Gallerani 2008; Dayal et al. 2009; Dayal, Ferrara \\& Saro 2010), that reproduces a number of important observational data sets accumulated for high-z LAEs. ",
        "conclusions": "\\label{discussion} We have linked the properties of high-$z$ LAEs to the Local Universe by coupling the semi-analytical code {\\tt GAMETE} to a previously developed LAE model. According to our results {\\it} the progenitors of MW-like galaxies cover a wide range of observed Ly$\\alpha$ luminosity, $L_{\\alpha} =10^{39-43.25}$ erg~s$^{-1}$, with $L_{\\alpha}$ increasing with $M_*$ (or, equivalently, $M_h$); hence some of them can be observed as LAEs. In each hierarchical merger history we find that, on average, {\\it only one} star-forming progenitor (among $\\approx 50$) is a LAE, usually corresponding to the major branch of the tree ($M_h\\approx 10^{10}\\Msun$). Nevertheless, the probability to have {\\it at least} one visible progenitor in any merger history is very high ($P=68\\%$). Interestingly, we found that the LAE candidates can be also observed as dropout galaxies since their UV magnitudes are always $M_{UV} < -18$.  On average the identified LAE stars have intermediate ages, $t_* \\approx 150-400$~Myr, and metallicities $Z\\approx 0.3-1 \\Zsun$; an exception is represented by five newly formed galaxies, which are hosted by small DM halos, $M_h\\approx 10^9\\Msun$, which are yet visible as (faint) LAEs ($L_{\\alpha}\\approx 10^{42.05}$~erg~s$^{-1}$) by virtue of their high star formation rate and extremely young stellar population, $t_* < 5$~Myr. The low metallicity of these young galaxies, $Z\\approx 0.016-0.044 \\Zsun$, reflects that of the MW environment at their formation epoch. Although rare (5 out of 80 LAEs), these Galactic building blocks could be even more unambiguously identified among the least luminous LAEs, due to their larger Ly$\\alpha$ equivalent widths, $EW = 60-130$~\\AA\\ with respect to those ($\\approx 40$~\\AA) shown by older LAEs\\footnote{Due to the low dust mass, $f_c= f_\\alpha \\approx1$ for all the progenitors we identify as LAEs; the observed EW is solely governed by $T_\\alpha$, which leads $EW \\approx 40$ \\AA\\, for almost all LAEs, except those with $t_*\\leq 10$ Myr.}. These small and recently virialized halos ($z_{vir}\\lsim 6$) could be the progenitors of Fornax-like dwarf spheroidal galaxies (see Fig.~1 of Salvadori \\& Ferrara 2009). By identifying these LAEs, therefore, it would be possible to observe the most massive dSphs of the MW system just at the time of their birth.  Uncertainties remain on the treatment of dust in calculating the LF and observed properties of the MW progenitor LAEs identified here, especially at the low luminosity end of the LF. Several aspects require additional study. As gas, metal and dust are preferentially lost from low mass halos, pushing the mass resolution of simulations to even lower masses would be important. Also, the amount of dust lost in SN-driven winds remains very poorly understood, an uncertainty that propagates in the evaluation of the continuum and Ly$\\alpha$ radiation escaping from the galaxy. Finally, Ly$\\alpha$ photons could be affected considerably by the level of dust clumping. It is unclear to what extent these effects may impact the visibility of LAEs, as discussed in e.g. Dayal, Ferrara \\& Saro (2010). Progress on these issues is expected when high-resolution FIR/sub-mm observations of LAEs with ALMA will become available in the near future (Finkelstein et al. 2009; Dayal, Hirashita \\& Ferrara 2010)."
    },
    "1005/1005.2562_arXiv.txt": {
        "abstract": "{Sensitive \\herschel\\ far-infrared observations can break degeneracies that were inherent to previous studies of star formation in high-z AGN hosts. Combining PACS 100 and 160$\\mu$m observations of the GOODS-N field with 2Msec \\chandra\\ data, we detect $\\sim$20\\% of X-ray AGN individually at $>3\\sigma$. The host far-infrared luminosity of AGN with $\\lhard\\approx 10^{43}\\ergs$ increases with redshift by an order of magnitude from z=0 to z$\\sim$1. In contrast, there is little dependence of far-infrared luminosity on AGN luminosity, for $\\lhard\\lesssim 10^{44}\\ergs$ AGN at z$\\gtrsim$1. We do not find a dependence of far-infrared luminosity on X-ray obscuring column, for our sample which is dominated by $\\lhard<10^{44}\\ergs$ AGN. In conjunction with properties of local and luminous high-z AGN, we interpret these results as reflecting the interplay between two paths of AGN/host coevolution. A correlation of AGN luminosity and host star formation is traced locally over a wide range of luminosities and also extends to luminous high z AGN. This correlation reflects an evolutionary connection, likely via merging. For lower AGN luminosities, star formation is similar to that in non-active massive galaxies and shows little dependence on AGN luminosity. The level of this secular, non-merger driven star formation increasingly dominates over the correlation at increasing redshift. } ",
        "introduction": "Measuring the star formation rate of the host galaxy is important for studying the co-evolution of active galactic nuclei (AGN) and their hosts. It is often difficult because the AGN can outshine the host at many wavelengths. However, the rest frame far-infrared/submm emission appears dominated by the host for AGN with $\\rm\\nu L_\\nu$(60$\\mu$m)$\\approx$0.1\\ldots 0.2 $\\rm L_{Bol,AGN}$ and higher (e.g. Netzer et al. \\cite{netzer07} and introduction to Lutz et al. \\cite{lutz10}, see also Wang et al. \\cite{wang08} for luminous high-z QSOs) and has been used as a host star formation rate diagnostic of high-z AGN (e.g. Serjeant \\& Hatziminaoglou \\cite{serjeant09}, Mullaney et al. \\cite{mullaney10}, and Lutz et al. \\cite{lutz10}). \\herschel\\ can detect much lower star formation rates than previous far-infrared and submm studies. It further improves such work by measuring the rest frame far-infrared SED peak without extrapolation from longer wavelengths. We here present a first \\herschel\\ study of rest frame far-infrared emission and host star formation in X-ray selected AGN in the GOODS-N field. Throughout the paper, we adopt an $\\Omega_m =0.3$, $\\Omega_\\Lambda =0.7$ and $H_0=70$ km\\,s$^{-1}$\\,Mpc$^{-1}$ cosmology. ",
        "conclusions": ""
    },
    "1005/1005.2612_arXiv.txt": {
        "abstract": "We compute  new chemical profiles for  the core and  envelope of white dwarfs  appropriate for pulsational  studies of  ZZ Ceti  stars. These profiles  are  extracted from  the  complete  evolution of  progenitor stars,  evolved through  the main  sequence and  the thermally-pulsing asymptotic giant branch (AGB)  stages, and from time-dependent element diffusion during  white dwarf evolution. We discuss  the importance of the initial-final mass relationship  for the white dwarf carbon-oxygen composition.  In particular, we find that the central oxygen abundance may be underestimated by about 15\\% if the white dwarf mass is assumed to be the hydrogen-free core  mass before the first thermal pulse.  We also discuss the importance for  the chemical profiles expected in the outermost  layers  of  ZZ  Ceti   stars  of  the  computation  of  the thermally-pulsing  AGB  phase  and  of  the  phase  in  which  element diffusion is relevant. We find  a strong dependence of the outer layer chemical stratification  on the stellar  mass.  In particular,  in the less massive models, the  double-layered structure in the helium layer built  up during  the thermally-pulsing  AGB phase  is not  removed by diffusion  by the  time the  ZZ Ceti  stage is  reached.   Finally, we perform adiabatic pulsation  calculations and discuss the implications of our new chemical profiles for the pulsational properties of ZZ Ceti stars.   We find  that  the  whole $g-$mode  period  spectrum and  the mode-trapping properties  of these  pulsating white dwarfs  as derived from our new chemical  profiles are substantially different from those based on chemical profiles widely used in existing asteroseismological studies. Thus,  we expect the asteroseismological  models derived from our chemical  profiles to be significantly different  from those found thus far. ",
        "introduction": "\\label{intro}  Pulsating DA (H-rich atmospheres) white dwarfs, also called ZZ Ceti or DAV stars, are  the most numerous class of  degenerate pulsators, with over  143 members  known  today  (Winget \\&  Kepler  2008). Since  the discovery of  the first ZZ  Ceti star, HL  Tau 76, by  Landolt (1968), there  has been a  continuous effort  to model  the interior  of these variable stars.  ZZ Ceti stars are found within a very narrow strip of effective  temperatures ($10\\,500$  K $\\lesssim  T_{\\rm  eff} \\lesssim 12\\,500$  K).   They  are  characterized by  multiperiodic  brightness variations  of  up to  $0.30$  mag  caused  by spheroidal,  non-radial $g$-modes of low  degree ($\\ell \\leq 2$) with  periods between 100 and 1200 s.  The  driving mechanism thought to excite  the pulsations near the  blue  edge  of  the  instability  strip  is  the  $\\kappa-\\gamma$ mechanism  that takes place  in the  hydrogen partial  ionization zone (Dolez \\&  Vauclair 1981; Dziembowski  \\& Koester 1981; Winget  et al. 1982).  Also,  the ``convective driving'' mechanism  has been proposed --- first by Brickhill (1991) and later re-examined by Goldreich \\& Wu (1999). It appears to be the  responsible of mode driving once a thick convection zone has developed at the stellar surface.  The comparison of  the observed pulsation periods in  white dwarfs and the periods  computed for appropriate theoretical  models (white dwarf asteroseismology) allows  to infer  details of their  origin, internal structure and  evolution (Winget \\& Kepler 2008;  Fontaine \\& Brassard 2008).  In  particular, the stellar  mass, the thickness of  the outer envelopes, the core chemical composition, magnetic fields and rotation rates can be  determined from the observed periods.   In addition, the asteroseismology  of ZZ  Ceti stars  is a  valuable tool  for studying axions (Isern  et al.  1992;  C\\'orsico et al.  2001;  Bischoff-Kim et al.  2008; Isern et al.   2010) and crystallization (Montgomery et al. 1999; C\\'orsico et  al.  2004, 2005; Metcalfe et  al.  2004; Kanaan et al.  2005). Finally, the temporal  changes in the observed periods can help  detect  planets  orbiting   around  white  dwarfs  (Mullally  et al. 2008).  The first published complete set of DA white dwarf models suitable for asteroseismology  was  that  of  Tassoul  et  al.   (1990).   A  large parameter space  was explored  in such a  monumental study, and  for a long time  (since the early  eighties) this set of  models represented the state-of-the-art  in the area.  The pulsation properties  of these models were  thoroughly explored  in a series  of important  papers by Brassard et  al.  (1991, 1992a,  1992b). As important as  these models were at  that time,  they suffer from  a number of  shortcomings.  For instance, the  core of  the models  is made of  pure  carbon,  while stellar evolution calculations indicate that cores of typical white dwarfs are made  of a  mixture of   carbon and oxygen.  Also,  the carbon/helium (C/He) and helium/hydrogen (He/H) chemical interfaces are modeled on the basis of the assumption of the diffusive equilibrium in  the ``trace element approximation'',  an approach that involves  a  quasi-discontinuity  in  the  chemical  profiles  at  the transition regions  which, in  turn, leads to  peaked features  in the Brunt-V\\\"ais\\\"al\\\"a  frequency and  exaggerated  mode-trapping effects (C\\'orsico  et al.   2002a, 2002b).   These models  were  employed for asteroseismological inferences of the DAVs G 226$-$29 (Fontaine et al. 1992) and GD 154 (Pfeiffer et  al.  1996).  More recently, Pech et al. (2006) and Pech \\&  Vauclair (2006) have presented asteroseismological analysis  on HL  Tau 76  and  G 185$-$32,  respectively, by  employing similar DA white dwarf models, although with updated input physics.  The models of Bradley  (1996) constituted a substantial improvement in the  field.    These  models  have  carbon-oxygen   cores  in  varying proportions,  and  the C/He  and  He/H  chemical  interfaces are  more realistic. Perhaps the most severe  shortcoming of these models is the (unrealistic)  ramp-like  shape  of  the core  carbon-oxygen  chemical profiles.   These DA  models  were  the basis  of  the very  important asteroseismological studies  on the DAVs G 29-38  (Bradley \\& Kleinman 1997), G  117$-$B15A and  R 548  (Bradley 1998), GD  165 and  L 19$-$2 Bradley (2001), and G 185$-$32 (Bradley 2006).  The next step in improving the modeling of DAVs was given by C\\'orsico et al. (2002b) and Benvenuto et al. (2002a), who employed evolutionary models  characterized by  He/H  chemical interfaces  resulting from  a time-dependent  element  diffusion  treatment  (Althaus  \\&  Benvenuto 2000), and  the carbon-oxygen  core chemical structure  extracted from the evolutionary computations  of Salaris et al.  (1997).   The use of very  smooth   outer  chemical  interfaces,  as   shaped  by  chemical diffusion, revealed  that the use  of the trace  element approximation turns  out to  be inappropriate  to model  the shape  of  the chemical interfaces in a  DA white dwarf.  This grid of  models was employed in an asteroseismological study of  G 117$-$B15A (Benvenuto et al. 2002a). For these  sequences, the starting configurations for  the white dwarf evolution were  obtained through an  artificial procedure, and  not as result of evolutionary computations of the progenitor stars.  Recently,  Castanheira \\&  Kepler  (2008, 2009)  have  carried out  an extensive  asteroseismological study  of  DAVs by  employing DA  white dwarf models similar to those of Bradley (1996), but with a simplified treatment  of the  core  chemical structure,  by somewhat  arbitrarily fixing  the central  abundances to  50 \\%  oxygen and  50 \\%  carbon.   The He/H chemical interfaces adopted for  these models are a parametrization of the realistic chemical  profiles resulting from time-dependent element diffusion (Althaus et al.  2003). The study includes the ``classical'' DAVs and also the recently discovered  SDSS DAVs. In total, 83 ZZ Ceti stars are analyzed.  An important  result of these studies is that the thickness of the H envelopes  inferred from asteroseismology is in the range $10^{-4}  \\gtrsim M_{\\rm H}/M_*  \\gtrsim 10^{-10}$, with  a mean value of  $M_{\\rm H}/M_*=  5 \\times 10^{-7}$.   This suggests  that an important  fraction of  DAs characterized  by  envelopes substantially thinner than  predicted by the standard evolution  theory could exist, with  the consequent important  implications for  the theory  of white dwarf formation.   However, these results  are preliminary and  do not include the  possible effects  of realistic carbon-oxygen  profiles on the asteroseismological fits.  Almost simultaneously  with the study of Castanheira  \\& Kepler (2008, 2009), Bischoff-Kim et al.  (2008) performed a new asteroseismological study on  G 117$-$B15A and  R 548 by  employing DA white  dwarf models similar to those  employed by Castanheira \\& Kepler  (2008, 2009), but incorporating realistic core chemical profiles according to Salaris et al. (1997).  The results of this work are in reasonable agreement with previous  studies on  these  ZZ  Ceti stars.   However,  the mass  and effective  temperatures found  by Bischoff-Kim  et al.   (2008)  for G 117$-$B15A  are  rather  high   (especially  the  mass,  at  $0.66  \\, M_{\\sun}$).  Recently, Bogn\\'ar et  al.  (2009) have employed the same asteroseismological modeling  to study the  pulsations of the  ZZ Ceti star  KUV  02464+3239.   Finally,  Bischoff-Kim (2009)  presented  the results  of an  asteroseismological  analysis of  two  DAVs with  rich pulsation spectrum, G 38$-$29 and  R 808 based on similar models, with parametrized, smooth  ramp-like core profiles.  These  models are able to reproduce the observed  period spectra reasonably well, though some assumptions  about the  $m$ and  $\\ell$ identification  of  modes were made.  White  dwarf  stellar  models  with realistic  chemical  profiles  are crucial  to  correctly  assess   the  adiabatic  period  spectrum  and mode-trapping properties of the DAVs,  which lies at the core of white dwarf  asteroseismology   (Brassard  et  al.    1992a;  Bradley  1996; C\\'orsico et  al.  2002a).   In this paper,  we compute for  the first time  consistent chemical  profiles for  both the  core {\\sl  and} the envelope of  white dwarfs with various stellar  masses appropriate for detailed asteroseismological  fits of  ZZ Ceti stars.   These chemical profiles  are  extracted  from  the  full and  complete  evolution  of progenitor   stars  from   the  zero   age  main   sequence,   to  the thermally-pulsing and mass-loss phases  on the asymptotic giant branch (AGB),  and from time-dependent  element diffusion  predictions during the white dwarf stage\\footnote{The  chemical profiles for the core and envelope of our models appropriate  for ZZ Ceti stars are available at {\\tt  http://www.fcaglp.unlp.edu.ar/evolgroup}}.  These  profiles will be  valuable in  conducting future  asteroseismological studies  of ZZ Ceti  stars  that  intend   to  include  realistic  chemical  profiles throughout the interior of white dwarfs. To assess the impact of these new  profiles on  the  theoretical pulsational  inferences, we  perfom adiabatic  pulsation computations  and compare  the  resulting periods with the pulsational inferences based on the most widely used chemical profiles in existing asteroseismological fits.  The  paper  is organized  as  follows.   In  Sect.  \\ref{physics},  we provide a description of the input physics assumed in the evolutionary calculations   of  relevance   for  the   chemical   composition.   In Sect. \\ref{mimf}  we discuss the importance of  the initial-final mass relationship for  the expected white  dwarf carbon-oxygen composition. The resulting chemical profiles are described at some length in Sect. \\ref{profiles}.  The implications of our new chemical profiles for the pulsational  properties  of  ZZ  Ceti  stars are  discussed  in  Sect. \\ref{pulsation}. We conclude in Sect. \\ref{conclusions} by summarizing our findings. ",
        "conclusions": "\\label{conclusions}  In  this paper  we computed  new chemical  profiles for  the  core and envelope  of white dwarfs  appropriate for  pulsational studies  of ZZ Ceti stars.   These profiles were  derived from the full  and complete evolution of progenitor stars from the zero age main sequence, through the  thermally-pulsing and  mass-loss phases  on the  asymptotic giant branch (AGB).  These new profiles are intented for asteroseismological studies  of ZZ  Ceti stars  that require  realistic  chemical profiles throughout  the white dwarf  interiors. In  deriving the  new chemical profiles,  we employed the  {\\tt LPCODE}  evolutionary code,  based on detailed  and  updated  constitutive physics.   Extra-mixing  episodes during  central hydrogen  and helium  burning,  time-dependent element diffusion during  the white dwarf stage  and chemical rehomogenization of   the   inner    carbon-oxygen   composition   by   Rayleigh-Taylor instabilities were considered. The  metallicity of progenitor stars is $Z=0.01$.  We discussed at  some length the importance of  the initial-final mass relationship  for  the   white  dwarf  carbon-oxygen  composition.   A reduction  of  the  efficiency  of extra-mixing  episodes  during  the thermally-pulsing  AGB   phase,  supported  by   different  pieces  of theoretical and  observational evidence, yields a  gradual increase of the hydrogen-free  core mass as evolution proceeds  during this phase. As a  result, the  initial-final mass relationship  by the end  of the thermally-pulsing AGB  is markedly different from  that resulting from considering the mass of the  hydrogen free core right before the first thermal  pulse.  We  found that  this issue  has implications  for the carbon-oxygen composition  expected in a white  dwarf.  In particular, the central oxygen abundance may be underestimated by about 15\\% if we assume the white  dwarf mass to be the  hydrogen-free core mass before the  first thermal  pulse.   We  also discuss  the  importance of  the computation of the thermally-pulsing AGB and element diffusion for the chemical profiles expected  in the outermost layers of  ZZ Ceti stars. In  this  sense, we  found  a strong  dependence  of  the outer  layer chemical stratification  on the stellar mass. In  less massive models, the intershell  region rich  in helium and  carbon --- which  is built during the  mixing episode at  the last AGB  thermal pulse ---  is not removed by diffusion by the time the ZZ Ceti stage is reached.  Finally, we  performed adiabatic pulsation  computations and discussed the implications  of our new  chemical profiles for  the pulsational properties  of ZZ  Ceti stars.   We  found that  the whole  $g-$mode period spectrum and the  mode-trapping properties of these pulsating white  dwarfs  as  derived   from  our  new  chemical  profiles  are substantially  different from those  based on  the most  widely used chemical  profiles  in   existing  asteroseismological  studies.  We expect the best fit parameters of asteroseismological studies using the {\\tt LPCODE} chemical profiles  to differ significantly from those found in studies made so far. Further studies will show in what way. Will we solve  the high mass  problem with  G117$-$B15A  and Salaris-like  core profiles (Bischoff-Kim et al. 2008) or find thicker hydrogen layers in asteroseismological  fits,   more  in  line   with  stellar  evolution calculations (Castanheira \\& Kepler 2008)?"
    },
    "1005/1005.0617_arXiv.txt": {
        "abstract": "We investigate the dark matter and the cosmological baryon asymmetry in a simple theory where baryon ($B$) and lepton ($L$) number  are local gauge symmetries that are spontaneously broken. In this model, the cold dark matter candidate is the lightest new field with baryon number and its stability is an automatic consequence of the gauge symmetry. Dark matter annihilation is either through a leptophobic gauge boson whose mass must be below a TeV or through the Higgs boson. Since the mass of the leptophobic gauge boson has to be below the TeV scale one finds that in the first scenario there is a lower bound on the elastic cross section of about $5 \\times 10^{-46} \\  \\rm{cm}^2$.  Even though baryon number is gauged and not spontaneously broken until the weak  scale, a cosmologically acceptable baryon excess is possible. There is tension between achieving both the measured baryon excess and the dark matter density. ",
        "introduction": "In the LHC era, we hope to either verify the standard model  or discover the theory that describes the physics of the weak scale. One of the open issues in the standard model (SM)  is the origin of the accidental global symmetries, $U(1)_{B}$ and $U(1)_{L}$, where $B$ stands for baryon number and $L$ for the total  lepton number. At the non-renormalizable level in the SM one can find operators that violate baryon number and  lepton number. For example, $QQQl/\\Lambda_{B}^2$ and $llHH/\\Lambda_L$, where $\\Lambda_B$ and $\\Lambda_L$ are the scales where $B$ and $L$ are respectively broken~\\cite{Weinberg:1979sa}.  Since the $QQQl/\\Lambda_{B}^2$ operator gives rise to proton decay~\\cite{Nath:2006ut} the cutoff of the theory has to be very large, $\\Lambda_{B} > 10^{15}$ GeV.  There is no other reason that the cutoff of the SM has to be that large and so it is worth thinking  about the possibility that both $B$ and $L$ are local gauge symmetries that are spontaneously broken~\\cite{FileviezPerez:2010gw} at a much lower scale (e.g., the weak scale)  and it is these gauge symmetries that prevent proton decay.  Recently, two simple models (denoted model (1) and model (2)) where $B$ and $L$ are local gauge symmetries have been proposed~\\cite{FileviezPerez:2010gw}.  In these models all anomalies are cancelled by adding a single new fermionic generation. One of the theories (model (1)) has an interesting realization of the seesaw mechanism~\\cite{Minkowski:1977sc, Gell-Mann1979, Mohapatra:1979ia} for neutrino masses and they both have a natural suppression of tree-level flavor changing neutral currents in the quark and leptonic sectors due to the gauge symmetries and particle content.  In model (2), the neutrinos have Dirac masses.  In addition, for model (2), the lightest new field with baryon number is a  candidate for the cold dark matter and its stability is an automatic consequence of the gauge symmetry.  It has been shown in Ref.~\\cite{FileviezPerez:2010gw} that $B$ and $L$ can be broken at the weak scale and one does not generate dangerous operators mediating proton decay.  We show how a dark matter candidate can arise in model (1).   In this article we investigate the properties of the cold dark matter candidates in the models proposed in Ref.~\\cite{FileviezPerez:2010gw} and study the implications of spontaneous $B$ and $L$ breaking at the weak scale for the baryon asymmetry in the Universe.  In model (2), the dark matter candidate, $X$, which has baryon number $-2/3$  can either annihilate through the leptophobic $Z_B$ present in the theory or through the Higgs boson.  We study the constraints from the relic density and the predictions for the elastic cross section relevant for direct detection experiments. We discuss the implications of the gauging of $B$ and $L$ for baryogenesis.  There is a potential conflict between the measured baryon excess and dark matter density.  For model (1), we discuss the generation of a baryon excess.  We introduce a limit of the theory where $L$ is broken at a high scale but $B$ is spontaneously broken at the weak scale.  In this limit  standard leptogenesis plus a primordial excess in the field responsible for baryon number breaking can give rise to an acceptable baryon excess and dark matter density even though the baryon number gauge symmetry is not broken until the weak  scale.  This paper is organized as follows: In Section \\ref{Section2} we discuss the main features of the model. In Section \\ref{Section3} we discuss, for model (2),  the properties of the dark matter candidate in the theory, constraints from the relic density and the predictions for the elastic cross section relevant for direct detection experiments. The properties of the dark matter candidate in model (1) are similar to cases already discussed in the literature (see for example \\cite{LopezHonorez:2006gr} and \\cite{Dolle:2009fn}). In Section \\ref{Section4} we discuss the implications of the breaking of $B$ and $L$ at the weak  scale for baryogenesis.  We summarize the main results in Section \\ref{Section5}. ",
        "conclusions": "\\label{Section5} We have investigated the cosmological aspects of two simple models, denoted (1) and (2), in which baryon number ($B$) and lepton number ($L$) are local gauge symmetries that are spontaneously broken around the weak scale.  In these models, the stability of our scalar dark matter candidate is a consequence of the gauge symmetry.  In model (2), we studied the possible dark matter annihilation channels and found what values of the masses and couplings lead to the observed relic abundance of dark matter.  In the case where the s-wave annihilation through an intermediate Higgs dominates, we find that, for $M_H = 120$ GeV, in order to evade the direct detection bounds the coupling between the Higgs and the dark matter must be less than $10^{-1.5}$ and $ 51 \\ \\rm{GeV}\\ \\lesssim M_X \\lesssim 63 \\  \\rm{GeV} $.  In the case where the p-wave annihilation through an intermediate leptophobic gauge boson dominates, we find that the coupling between the leptophobic $Z_B$ and the dark matter must be less than $0.1$ and $ 235 \\ \\rm{GeV}\\ \\lesssim M_X \\lesssim 250 \\  \\rm{GeV} $ when $M_{Z_B}=500$ GeV.  In this case the leptophobic gauge boson has to be below the TeV scale and one finds a lower bound on the elastic cross section $\\sigma_{SI}^B \\gtrsim 5 \\times 10^{-46} \\  \\rm{cm}^2$. In both cases, direct detection experiments constrain the annihilation to proceed close to resonance in order to evade direct detection and to produce the observed relic abundance of dark matter. We have shown that even though baryon number is gauged and spontaneously broken at the weak scale it is possible to generate a cosmological baryon excess.  A modest fine-tuning is needed to achieve both the measured dark matter relic abundance and baryon excess.  In model (1), we introduced a simple mechanism to split the masses of the real of the imaginary part of the neutral component of the new scalar doublet to evade direct detection limits.  We showed that one can simultaneously achieve both the observed baryon asymmetry of the Universe and the dark matter relic abundance.  In particular,  when $L$ is broken at the high scale but $B$ is spontaneously broken at the weak scale, standard leptogenesis can be applied.   \\subsection*"
    },
    "1005/1005.4728_arXiv.txt": {
        "abstract": "It is found that \\feii emission contributes significantly to the optical and ultraviolet spectra of most active galactic nuclei. The origin of the optical/UV \\feii emission is still a question open to debate. The variability of \\feii would give clues to this origin. Using 7.5 yr spectroscopic monitoring data of one Palomer-Green (PG) quasi-stellar object (QSO), PG 1700+518, with strong optical \\feii emission, we obtain the light curves of the continuum \\lv, \\feii, the broad component of \\hb, and the narrow component of \\hb by the spectral decomposition. Through the interpolation cross-correlation method, we calculate the time lags for light curves of \\feii, the total \\hb, the broad component of \\hb, and the narrow component of \\hb with respect to the continuum light curve. We find that the \\feii time lag in PG1700+518 is $209^{+100}_{-147}$ days, and the \\hb time lag cannot be determined. Assuming that \\feii and \\hb emission regions follow the virial relation between the time lag and the FWHM for the \\hb and \\feii emission lines, we can derive that the \\hb time lag is $148^{+72}_{-104}$ days. The \\hb time lag calculated from the empirical luminosity--size relation is 222 days, which is consistent with our measured \\feii time lag. Considering the optical \\feii contribution, PG 1700+518 shares the same characteristic on the spectral slope variability as other 15 PG QSOs in our previous work, i.e., harder spectrum during brighter phase. ",
        "introduction": "The variability is a common phenomenon in quasi-stellar objects (QSOs) and provides a powerful constraint on their central engines. In the past two decades, the optical variability research focused on the spectral monitoring instead of the pure photometric monitoring. With the active galactic nuclei (AGNs) watch and the Palomer-Green (PG) QSOs spectrophotometrical monitoring projects, the reverberation mapping method, i.e., exploring the correlation between the emission lines and the continuum variations, is used to investigate the inner structure in AGNs \\citep[e.g.,][]{Blandford82, Peterson93}. It is found that motions of clouds in the broad line regions (BLRs) are virialized \\citep[e.g.,][]{Kaspi00, Kaspi05, Peterson04}.  With the line width of \\hb, \\mgii, \\civ from BLRs, the empirical size-luminosity relation derived from the mapping method is used to calculate the masses of their central supermassive black holes \\citep[SMBHs; e.g.,][]{Kaspi00, McLure04, Bian04, Peterson04, Greene05}.  It is found that the \\feii emission contributes significantly to the optical and ultraviolet spectra of most AGNs. Thousands of UV \\feii emission lines blend together to form a pseudocontinuum, resulting in the ``small blue bump'' around 3000\\AA\\ when they are combined with Balmer continuum emission \\citep[e.g.,][]{Wills85}. The optical \\feii would lead to two bumps in two sides around the \\hb $\\lambda 4861$\\AA\\ \\citep[e.g.,][]{Boroson92}. It is found that the flux ratio of \\feii to \\hb, $R_{\\rm Fe}$, where the optical \\feii flux is the flux of the \\feii\\ emission between $\\lambda$4434 and $\\lambda$4684, strongly correlates with the so-called Eigenvector 1, which is suggested to be driven by the accretion rate \\citep[e.g.,][]{Boroson92, Marziani03a}.  The origin of the optical/UV \\feii emission is still an open question. It is found that photoionized BLRs cannot produce the observed shape and strength of the optical \\feii emission and that the strength of UV \\feii cannot be explained unless considering the micro-turbulence of hundreds of \\kms or the collisional excitation in warm, dense gas \\citep{Baldwin04}. However, \\citet{Vestergaard05} found the correlation between the optical \\feii variance and the continuum variance and suggested that the optical \\feii is due to the line fluorescent in a photoionized plasma. It suggests that the optical \\feii line do not come from the same region as the UV \\feii emission \\citep[e.g.,][]{Kuehn08}. \\citet{Maoz93} found that the reverberation time lag of UV \\feii in NGC 5548 is about 10 days, similar to \\civ time lag, smaller than the \\hb time lag. The reverberation measurement for the optical \\feii emission has not fared so well. Some suggested that the optical emission is produced in the same region as the other broad emission lines, and some suggested that it is in the outer portion of the BLRs because of narrower FWHM of \\feii with respect to \\hb \\citep[e.g.,][]{Laor97, Marziani03a, Vestergaard05, Kuehn08}. Recently, \\citet{Hu2008a, Hu2008b} did a systematic analysis of \\feii emission in QSOs from the Sloan Digital Sky Survey (SDSS), and found that the \\feii emission is redshifted with respect to the rest frame defined by the \\oiii narrow emission line and \\hb intermediate-width component is correlated with \\feii which locates at the outer portion of the BLRs.  \\citet{Kaspi00} gave the 7.5 yr spectroscopic monitoring data for 17 PG QSOs. There is one PG QSO, PG 1700+518, with strongest optical \\feii emission and $R_{\\rm Fe}=1.42$ \\citep{Turnshek85, Boroson92}. Its \\hb FWHM is $1846\\pm 682 \\kms$ \\citep{Peterson04}, and it is also called as a narrow line Seyfert 1 galaxy (NLS1). Using the \\feii template from one NLS1, I ZW 1, we model the \\feii emission to investigate the \\feii variability and the relation to the continuum variability in PG 1700+518. The data and analysis are described in Section 2, the results are given in Section 3, the discussion is given in Section 4, and the conclusions are presented in Section 5. All of the cosmological calculations in this paper assume $H_{\\rm 0}=70 \\kms \\rm ~Mpc^{-1}$, $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "With the spectral decomposition of 39 spectra of PG 1700+518 with the strong \\feii emission, we investigate the \\feii variability and its time lag. The main conclusions can be summarized as follows: (1) we give light curves of \\lv, \\feii, \\hbb, \\hbn, and \\hb$^{n+b}$, as well as the mean and rms spectra for PG 1700+518. With the normalized variability measure, $\\sigma_N$, we find that all components are variable. (2) With the code of Peterson et al. (2004), we find that \\feii time lag in PG1700+518 is $209^{+100}_{-147}$ days, and \\hb time lag cannot be determined. (3) Considering the uncertainties of time lags, the expected \\hb time lag from the empirical luminosity--size relation is 221.6 lt-days, consistent with our measured \\feii time lag. If we take FWHM/time lag uncertainties into consideration, \\feii emission region is located near the \\hb emission region, not conclusively located outside of the \\hb emission region. (4) Assuming that \\feii and \\hb emission regions follow the virial relation between the time lag and the FWHM for the \\hb and \\feii emission lines, we can derive that the \\hb time lag is $148^{+72}_{-104}$ days. With respect to the \\hb time lag of 252 days suggested by Peterson et al. (2004), smaller \\hb time lag, which leads to the black hole mass estimation in the logarithm, is decreased by 0.23 dex. (5) After considering \\feii contribution, PG 1700+518 shares the same characteristic on spectral slope variability to other 15 PG QSOs in our previous work \\citep{Pu06}, i.e., harder spectrum during brighter phase."
    },
    "1005/1005.3422_arXiv.txt": {
        "abstract": "{We present an analysis of the first space-based far-IR-submm observations of M~33, which measure the emission from the cool dust and resolve the giant molecular cloud complexes.  With roughly half-solar abundances, M33 is a first step towards young low-metallicity galaxies where the submm may be able to provide an alternative to CO mapping to measure their H$_2$ content.  In this Letter, we measure the dust emission cross-section $\\sigma$ using SPIRE and recent CO and \\HI\\ observations; a variation in $\\sigma$ is present from a near-solar neighborhood cross-section to about half-solar with the maximum being south of the nucleus.  Calculating the total H column density from the measured dust temperature and cross-section, and then subtracting the \\HI\\ column, yields a morphology similar to that observed in CO.  The H$_2$/\\HI\\  mass ratio decreases from about unity to well below 10\\% and is about 15\\% averaged over the optical disk.  The single most important observation to reduce the potentially large systematic errors is to complete the CO mapping of M~33. } ",
        "introduction": "Understanding star formation requires studying the interplay between the phases of the interstellar medium (ISM).  Dust processes most of the energy transiting the ISM, but the cool dust component, although representing the vast majority of the dust mass, is difficult to observe from the ground. {\\it Herschel} SPIRE observations \\citep{Pilbratt10,Griffin10} are the first space-based 250-500 \\micron\\ data and as such provide a unique occasion to put together a global picture of the cool gas and dust in M~33.  In particular, we compare the morphology of the Far-IR emission and that of the gas as determined from CO and \\HI\\ measurements and attempt to measure how the dust cross-section varies in M~33. A longer term goal is to be able to use the dust emission to constrain the variation of the $\\ratioo$ factor within M~33 and elsewhere.  This Letter is one of a series on the {\\tt HERM33ES} project on the ISM of the Local Group galaxy M~33, an overview of which is given in \\citet{Kramer10}, hereafter K10. For consistency with the other M\\,33 papers in this volume, we adopt a distance of $D$ = 840 kpc for M\\,33 (i.e., 25 arcsec = 100 pc) and orientation parameters of $PA$ = 22.5 degrees and $i$ = 56\\degr. We use the recent  {\\tt HERM33ES} SPIRE observations at 250, 350, and 500 \\micron\\, combined with CO(2--1) observations from the M33CO@IRAM project, as a tracer of the molecular component, and a high-resolution mosaic of VLA HI data \\citep[both from][]{Gratier10}. The SPIRE data were first processed as described in K10 and then converted from Jy/beam to brightness units (MJy/sr). To estimate the dust temperature, the 250\\micron\\ data were convolved to the 350\\micron\\ beamsize ($\\sim 25\"$) and, assuming a single temperature grey body with an emissivity $\\propto \\nu^\\beta$ with $\\beta =2$, a temperature was derived from the flux ratio. At the temperatures of the cool component seen in M~33, \\citet{Dupac03} find $\\beta \\approx 2$. To minimize the effect of the uncertainty in $\\beta$, we chose to use adjacent bands to estimate temperature. The 250/350\\micron\\ ratio provides more accurate temperatures than the 350/500 micron ratio even for temperatures below 10~K.  A 15\\% variation (or uncertainty) in the 250/350\\micron\\ ratios corresponds to a temperature change of 0.9, 2.1, and 4.1 K at temperatures of 10, 15, and 20 K but the same uncertainty in the 350/500\\micron\\ ratios yields 1.3, 3.2, and 6.3 K errors for the same dust temperatures.  A further advantage is that we obtain the dust temperature at a resolution typical of giant molecular clouds (GMC), $\\sim 100$ pc.  Ideally, a multi-component fit would be used but this requires high S/N data at many wavelengths over the whole disk.  By assuming that the cool dust component dominates the emission beyond 250\\micron, we could calculate the dust temperature out to the optical radius of M~33.  The resulting temperature map is shown in Fig. 1.  Figure 2 shows a comparison with the temperature of the cool component of the preferred two-component (warm plus cool dust) model fit to data between 24\\micron\\ and 500\\micron\\ from K10. The two-component model uses $\\beta=1.5$ and the temperatures are higher out to 6~kpc. The temperatures in the radial bins have been estimated by averaging the temperatures rather than averaging the emission as in K10.  The latter yields slightly higher temperatures because the dust is usually warmer where emission is strong (compare Fig. 1 with Fig 1. in K10). K10 calculate the dust column density using $\\kappa=0.4 (\\nu/250{\\rm GHz})^2 \\ {\\rm cm}^2$g$^{-1}$ of dust, which is equivalent to a dust cross-section per H-atom of $\\sigma = 1.1 \\times 10^{-25} (\\lambda/250\\mu {\\rm m})^{-2} {\\rm cm}^2$ for a hydrogen gas-to-dust mass ratio of 140. In the following we use the SPIRE 250/350 \\micron\\ color temperature because the data cover a greater area and agree well with the cool dust temperature of the two-component fit, showing the domination of the cool dust at these wavelengths. No correction for line contamination was subtracted from the SPIRE data.  At these frequencies, the CO lines contribute very little to the continuum flux, unlike at 1.3mm or 850\\micron\\ \\citep[e.g.][]{Braine_n3079}.  Figure 3 shows the SPIRE 250\\micron\\ emission with the CO(2--1) emission as contours at 25$''$ resolution to show how closely the CO emission follows the dust emission peaks. \\citet{Gratier10} show a detailed comparison of the CO and \\HI\\ emission with star formation tracers such as H$\\alpha$ and {\\it Spitzer} 8 and 24\\micron\\ maps. The  250\\micron -bright regions are detected in CO and the general morphology of the cool dust emission is seen in the \\HI\\ image.  \\begin{figure}[h] \\centering \\includegraphics[width=6.2cm,angle=-90]{tdust_rad4b.eps} \\\\ \\caption{Radial distribution of the average dust temperature in the stellar disk of M\\,33. The solid line joining the triangles shows the 250/350 \\micron\\ color temperature and for comparison we show the cool dust temperature (large dots) obtained by K10 in their preferred two-component model with $\\beta=1.5$. The error bars on the triangles indicate a 15\\% calibration uncertainty in the 250/350 \\micron\\ ratio converted to temperature.} \\end{figure}  \\begin{figure}[h] \\centering \\includegraphics[width=8.8cm,angle=0]{s250_co_figb.eps} \\\\ \\caption{SPIRE 250\\micron\\ emission with the IRAM CO(2--1) emission at 25$''$ resolution superposed as contours and with the polygons to the upper left used to calculate the dust cross-section per H-atom. The  beams are shown as white dots in the lower left corner.  The CO contours are at 0.5 (red), 2 (blue), and 4 K km s$^{-1}$ in the main beam scale.  The large polygon indicates the area covered so far by the IRAM CO survey.} \\end{figure} ",
        "conclusions": ""
    },
    "1005/1005.3278_arXiv.txt": {
        "abstract": "We report CCD photometry of the cataclysmic variable V1113 Cygni. During two campaigns, lasting from May to August 2003 and from March to June 2005, we recorded two superoutburst. In the obtained light curves we detected clear superhumps with a mean period  $P_{\\rm sh}=0.07891(3)$ days ($113.63\\pm0.04$ min). That fact confirms that the star is a member of SU UMa class of dwarf novae. During the first observed superoutburst the superhump period was decreasing with an enormous rate of $\\dot P = -4.5(8)\\times 10^{-4}$ which is one of the highest values ever observed in SU UMa systems.  \\noindent {\\bf Key words:} \\textit{Stars: individual: V1113 Cyg - binaries: close - novae, cataclysmic variables} ",
        "introduction": "Cataclysmic variable stars are close binary systems containing white dwarf (the primary) and red dwarf (the secondary). In those systems the secondary loses its mass through the inner Lagrangian point in the Roche-lobe and the primary star accretes it (Warner 1995).  SU UMa class of dwarf novae is one of the subclasses of cataclysmic variable stars. SU UMa-type stars show two types of outbursts: short (outbursts) and long (superoutbursts). Superoutbursts are about one magnitude brighter than ordinary outbursts and have a duration a few days longer than the outbursts. In the light curves of superoutbursts one can see periodic light oscillations (superhumps) with a period a few percent longer than the orbital period of the binary system. The thermal-tidal instability model proposed by Osaki (1996) is now known as the best explanation of the mechanism of the superoutbursts and superhumps (for a review see Warner 1995, Osaki 1996, Hellier 2001). However, in recent years some serious criticism concerning both the nature of superoutburst and nature of superhumps has appeared (Schreiber and Lasota 2007, Smak 2004, 2008, 2009).  V1113 Cyg was discovered as a dwarf nova by Hoffmeister (1966), but received little attention for almost 30 years. Results of Kato et al. (1995) have shown that the star is SU UMa-type with superhumps' period of $0.0792(\\pm0.0001)$ d. They also reported that V1113 Cyg has a large ($\\sim$ 6 mag) superoutburst amplitude.  In the second paper by Kato et al. (2001) the detection of 30 outburst from July 1994 to May 2001 is presented. They noticed low number of outbursts (only about two normal outbursts per each supercycle) and high number of superoutbursts (12 from 30 detected). They suspected variation of the mass transfer rate. Because results are inconclusive further observations were needed.  Kato et al. (2009) reanalyzed the observations of Kato et al. (1995) and obtained new observations during the 2008 superoutburst. They reported decreasing superhump periods in both superoutburts with the rates of $-19.2(6.8)\\times 10^{-5}$ and $-5.2(4.7)\\times 10^{-5}$, respectively. ",
        "conclusions": "Recently, Kato et al. (2009) published a comprehensive survey of period variations of superhumps in SU UMa-type dwarf novae. The conlcuded that in many systems evolution of the superhump period can be divided into three distinct parts: an early stage with a stable and longer period (stage A), middle stage for which many systems (especially with $P_{\\rm sh}<0.08$ days) show positive superhump period derivative (stage B) and final stage with a shorter but again stable superhump period (stage C). However they reported that V1113 Cyg does not follow this scenario with decreasing superhump periods with the rates of $-19.2(6.8)\\times 10^{-5}$ and $-5.2(4.7)\\times 10^{-5}$ observed for 1994 and 2008 superoutbursts, respectively.  V1113 Cyg has its superhump period close to 0.08 day borderline and thus can show different superhump period evolution. Our high negative $\\dot P$ value may indicate that the behaviour of the star is similar to that observed in SU UMa or DM Dra (see Figs. 4 and 7 in Kato et al. 2009).  On the other hand, it is also possible that in 2003 our observations caught V1113 Cyg during the transition from stage A to B. As can be estimated from Fig. 4 of Kato et al. (2009), for stars with superhump period of  around 0.078 days, such a transition occurs around cycle number $E\\approx 20$ assuming $E=0$  for the begining of the superoutburst (see $O-C$ diagrams of TT Boo, QY Per, RZ Leo or EG Aqr). The peak of the parabola shown in our Fig. 5 occurs also at $E\\approx 20$ but in this case the cycle numeration starts with the first night of the observations which is not the first night of the superoutburst. Assuming that the superoutburst started 1-2 days before our first observation the transition moment shifts to $E\\approx 32-33$ or even $E\\approx 45$. It is quite late in comparison to other stars putting some doubt in A to B transition hypothesis.  To summarize the results of the observations of the 2003 and 2005 brightenings of V1113 Cyg we can confirm: \\begin{itemize} \\item Detection of the clear superhumps during those events directly proves that V1113 Cyg belongs to the group of SU UMa variables. \\item The amplitude of the superoutburst is at least ($\\sim$ 4) mag, but we never caught the star at the beginning of the superoutburst. According to Kato (2001) this variable has a large ($\\sim$ 6) mag outburst amplitude and it is a possible estimation. \\item During two years of observation the star went into at least two superoutbursts so V1113 Cyg has about one superoutburst a year. Eruptions of the star lasted about two weeks and were not shorter than 12 days (July 2005 superoutburst). \\item We did not record any normal outbursts. Lack of  them may be an observational effect. But in general we agree with the results of Kato (2001) which suggest low ratio of outbursts to superoutbursts. \\item During the first observed superoutburst the superhump period was decreasing with an enormous rate of $\\dot P = -4.5(8)\\times 10^{-4}$ which is one of the highest values ever observed in SU UMa systems. \\end{itemize}  \\bigskip \\noindent {\\bf Acknowledgments.} ~We acknowledge generous allocation of the Warsaw Observatory 0.6-m telescope time.  Data from AAVSO observers are also appreciated. We thank to W. Pych for providing some useful software which was used in the analysis and to K. Mularczyk and P. K\\c{e}dzierski for their assistance in observations. This work was supported by the Polish MNiSW grant no. N203~301~335. KZ was supported by Foundation for the Polish Science through grant MISTRZ.  {\\small"
    },
    "1005/1005.3570_arXiv.txt": {
        "abstract": "On 1998 November 14, Saturn and its rings occulted the star GSC 0622-00345.  The occultation latitude was 55{\\decdegree}5 S.  This paper analyzes the \\mbox{2.3 {\\micron}} light curve derived by \\citeauthor{HarringtonFrench2010apjsatoc98I}.  A fixed-baseline isothermal fit to the light curve has a temperature of \\mbox{140 {\\pm} 3 K}, assuming a mean molecular mass of 2.35 AMU.  The thermal profile obtained by numerical inversion is valid between 1 and 60 {\\micro}bar.  The vertical temperature gradient is \\math{>}0.2 K km\\sp{-1} more stable than the adiabatic lapse rate, but it still shows the alternating-rounded-spiked features seen in many temperature gradient profiles from other atmospheric occultations and usually attributed to breaking gravity (buoyancy) waves.  We conduct a wavelet analysis of the thermal profile, and show that, even with our low level of noise, scintillation due to turbulence in Earth's atmosphere can produce large temperature swings in light-curve inversions.  Spurious periodic features in the ``reliable'' region of a wavelet amplitude spectrum can exceed 0.3 K in our data.  We also show that gravity-wave model fits to noisy isothermal light curves can lead to convincing wave ``detections''.  We provide new significance tests for localized wavelet amplitudes, wave model fits, and global power spectra of inverted occultation light curves by assessing the effects of pre- and post-occultation noise on these parameters.  Based on these tests, we detect several significant ridges and isolated peaks in wavelet amplitude, to which we fit a gravity wave model.  We also strongly detect the global power spectrum of thermal fluctuations in Saturn's atmosphere, which resembles the ``universal'' (modified Desaubies) curve associated with saturated spectra of propagating gravity waves on Earth and Jupiter.  \\if\\submitms y \\else \\comment{\\hfill\\herenote{DRAFT of {\\today} \\now}.} \\fi ",
        "introduction": "\\label{intro}  Earth-based occultations remain an attractive method for measuring the thermal profile in the 1--100-{\\micro}bar region of a planetary atmosphere.  Many such profiles for Saturn were recorded during the 28 Sgr occultation of 1989 July 3, which sampled the equatorial region from 6{\\decdegree}6 N--15{\\decdegree}2 S latitude \\citep{HubbardEtal1997icsatmes}.  There is a single profile for the north polar region \\citep[82{\\decdegree}5--85{\\degrees} N]{CoorayEtal1998icnpsatoc}, and a northern low-latitude profile from the same event \\citep[19{\\decdegree} 16 N]{FrenchEtal1999phemu97satoc}. Saturn's central flash probes much deeper, around 2.5 mbar; \\citet{NicholsonEtal1995icsatcfl} obtained IR images of the flash during the 28 Sgr event, from which they inferred the zonal wind profile of the sampled latitudes along Saturn's limb.  Occultations observed by a spacecraft near a giant planet probe the troposphere from the cloud deck (\\sim1 bar) to the mbar level at radio and infrared wavelengths.  They probe the upper stratosphere and thermosphere (\\math{<}1 {\\micro}bar) in the ultraviolet.  Earth-based visual and infrared occultations measure the thermal structure of the intervening mesosphere and stratosphere regions, which are not well sampled by spacecraft experiments.  For Saturn, the {\\em Pioneer} radio \\citep{KlioreEtal1980scipioneersatoc, LindalEtal1985ajvoysatoc} and {\\em Voyager 2} extreme ultraviolet solar and stellar observations \\citep{SmithEtal1983jgrvoysatoc} sensed the equatorial region only. The {\\em Voyager 1} radio occultation sensed 75{\\degrees} S \\citep{TylerEtal1981scisatvoy1radoc}, while the {\\em Voyager 2} radio occultations sensed 36{\\decdegree}5 N and 31{\\degrees} S \\citep{TylerEtal1982scivoy2radoc}.  The {\\em Cassini} radio experiment has performed a number of radio occultations in the equatorial region as well as at middle and high latitudes of Saturn \\citep{NagyEtal2006jgrCassinisatradio, KlioreEtal2009jgrSatOc}.  The {\\em Cassini} Ultraviolet Imaging Spectrograph stellar occultation probed the upper atmosphere at 40{\\degrees} S and 66{\\degrees} N \\citep{Shemansky2008cosparSatcassiniuvis}, and the {\\em Cassini} Composite Infrared Spectrometer mapped Saturn's thermal atmospheric emission, resulting in temperature maps for both hemispheres for pressures ranging from 0.1 mbar to about 700 mbar \\citep{FlasarEtal2005sciSatcassini}.  The {\\em Cassini} Visual and Infrared Mapping Spectrometer also has the capability of observing spectrally resolved near-infrared stellar occultations by Saturn's atmosphere \\citep{BrownEtal2004ssrVIMS}.  Temperature profiles for a variety of atmospheres from both occultations and in situ observations show quasi-periodic structures that are usually attributed to propagating waves.  Waves have been reported on Venus \\citep{HinsonJenkins1995icvenusradoc}, Earth (\\citealp{FrittsAlexander2003revgeogwdyn}, and references therein), Mars \\citep{CreaseyEtal2006grlMarsgw, FrittsEtal2006jgrMarsgwaero}, Jupiter \\citep{FrenchGierasch1974jasocwav, YoungEtal1997scijupgravwav, RaynaudEtal2003icjupoc, RaynaudEtal2004icjupbsco, YoungEtal2005icjupasi}, Saturn \\citep{CoorayEtal1998icnpsatoc, FouchetEtal2008natSaturn}, Titan \\citep{SicardyEtal1999ictit28sgr}, Uranus \\citep{YoungEtal2001icuoc}, Neptune \\citep{RoquesEtal1999aanepstrat3}, and Pluto \\citep{PersonEtal2008ajWavPluto, HubbardEtal2009icPlutowav, ToigoEtal2010icPlutoTides}.  Both the behavior of individual waves and the form of wave power spectra can reveal properties of the underlying atmosphere.  For example, the forcing, propagation, and dissipation of the waves both contribute to and depend on the sources and sinks of energy in the atmosphere, the background thermal state, and eddy and molecular diffusion.  On 1998 November 14, Saturn and its rings occulted GSC 0622-00345, as predicted by \\citet{BoshMcDonald1992ajsatocs}.  We obtained a light curve for atmospheric immersion, based on infrared imaging observations at the NASA Infrared Telescope Facility (IRTF) on Mauna Kea, HI.  The high signal-to-noise ratio (S/N) allowed us to determine the vertical temperature profile of Saturn's stratosphere at 55{\\decdegree}5 S latitude, a region not sampled by previous stellar occultation observations (see Figure 1 and Table 1 of \\citealp{HarringtonFrench2010apjsatoc98I}, hereafter Paper I).  Paper I presents the light curve and describes the new methods used to acquire and derive it.  This paper presents the scientific analysis of the light curve.  Subsequent sections cover isothermal model fits, numerical inversions to derive the thermal profile, noise tests, local and global wavelet spectrum analysis, a gravity-wave model based on wavelet reconstruction, exploration of the ``universal'' power spectrum of gravity waves, discussion of the global power spectrum, and our conclusions.  For each analysis, we present new significance tests that determine the effects of real (non-Gaussian) noise. ",
        "conclusions": "We have analyzed a light curve based on IRTF observations of the 1998 November 14 occultation of GSC 0622-00345 by Saturn (Paper I).  We presented and analyzed isothermal light-curve fits, an atmospheric thermal profile, a wavelet analysis, gravity wave modeling, and power spectra.  The derived thermal profile varies over 142 -- 151 K in the pressure range 1 -- 60 {\\micro}bar at a latitude of 55{\\decdegree}5 S. The vertical temperature gradient is removed by more than 0.2 K km\\sp{-1} from the adiabatic lapse rate, indicating that the stratospheric region sounded by the occultation is statically stable. Our thermal gradient profile shows the same alternating-rounded-spiked appearance of other occultation profiles, including one for Saturn \\citep{CoorayEtal1998icnpsatoc}.  This shape has previously been interpreted as evidence of gravity wave breaking \\citep{RaynaudEtal2003icjupoc, RaynaudEtal2004icjupbsco, YoungEtal2005icjupasi}.  Our new noise tests, based on real noise sampled from our light curve's upper baseline rather than synthetic, uncorrelated Gaussian noise, showed that atmospheric scintillation (and similar correlated noise sources such as spacecraft pointing drifts) can introduce relatively strong, spurious temperature fluctuations into the thermal profile derived by inverting an atmospheric occultation light curve.  For our data, the effect was many K for the raw inversions, but it mainly appeared at longer wavelengths in power spectra.  However, the amplitudes at the shorter wavelengths found outside the COI were still of order 0.1 -- 1 K, comparable to the amplitudes of the gravity waves often identified in occultation data sets.  We thus developed several significance tests for our power-spectrum analyses, and note that without such tests, one must be skeptical of gravity-wave detection claims in ground-based (and possibly some space-based) occultation inversions.  We used a wavelet analysis to search for localized gravity wave trains.  Based on the wavelet power and the use of significance tests only, the strongest candidate had an amplitude of 0.7 K, \\math{\\lambda\\sb{z}} = 12.5 km at \\math{z\\sb{\\rm max} = -35 km}, and lasted four cycles.  It and several shorter-wavelength features stand over the global power spectrum at the 95% Alternative explanations for the observed periodic structures include sound waves, planetary waves, and non-transient features.  Without knowing the horizontal structure of the wave we cannot rule out sound waves or planetary waves as the cause for the temperature fluctuations.  However, one can make the argument that, for a given \\math{\\lambda\\sb{z}}, planetary waves typically have lower (time) frequencies and therefore dissipate lower in the atmosphere than the valid region of our profile.  To take into account both phase and amplitude information, which could improve sensitivity over our amplitude-based noise limit, we fit a gravity-wave model to the amplitude and phase of the strongest features in the valid region of the wavelet spectrum, following \\citet{RaynaudEtal2004icjupbsco}.  The model calculates the temperature amplitude vs.\\ height of a single, damped wave mode propagating throughout our profile's valid region.  We performed fits both with and without a new parameter for constant vertical wind shear.  The added parameter did not improve the fits, so we report the shearless fits here for consistency with prior results from this model.  Our best candidate for a gravity wave that propagates continuously through the valid region, as assessed by our model, had \\math{\\lambda\\sb{z} = 40} km, \\math{\\lambda\\sb{h} = 200} km, and a period of 42 minutes.  At this altitude the wave is not strongly affected by dissipation and achieves a maximum amplitude above the observed atmospheric region exceeding 1 K.  According to Equation (\\ref{eq:detectability}), this wave was detectable if the angle between the line of sight and the horizontal direction of wave propagation exceeded 74{\\degrees}.  However, similar fits to quasi-periodic features in thermal profiles derived from isothermal (i.e., waveless) light curves with real noise gave some fits of similar appearance.  This demonstrates the need for a study of real noise to establish an amplitude criterion that discriminates real waves from noise.  The reason wavelike (i.e., sinusoidal) features arise out of the noise is simple: by reconstructing only a limited range of wavelengths, the resulting profile is certainly sinusoidal and has a favored period.  Wave models are sinusoidal, and ours even has parameters that allow the phase and amplitude to vary, so we will get a good fit if the amplitude of the reconstructed data does not vary much.  We find that this circumstance occurs in the noise data sets often enough to require at least criteria for significant wave amplitude.  Criteria involving the number of cycles outside the COI or the phase of the reconstructed data could potentially provide even-more-stringent limits.  We derived global power spectra from our wavelet transforms, using only data outside the COI.  This method greatly reduced the noise level of the spectra, which stand everywhere 2 -- 10 times above the noise level calculated using real-noise-contaminated isothermal light curves.  The power spectra follow the modified-Desaubies form of the universal spectrum of gravity waves, though with a slightly more negative high-wavenumber exponent.  Superposed on this spectrum one sees the signature of the discrete wave structures discussed above. That we see both the universal spectrum and individual features fit well by gravity-wave models lends confidence that we are indeed looking at a signal dominated by gravity waves.  The amplitudes of all five wave-like features that we analyzed are well above the noise level.  The wave model used to fit these features is based on the assumption that the waves propagate independently of each other.  Interactions between the wave modes might be able to explain the discrepancies between the observed wave fluctuations and the model's single gravity waves.  Note that the discrepancies are more significant when derived wavenumbers are larger than the observationally derived characteristic wavenumber.  This is the saturated part of the spectrum where wave-wave interactions determine wave behavior and vertical propagation.  Direct comparison of our derived temperatures with previous Saturn occultation measurements (see, e.g., Table VII of \\citealp{HubbardEtal1997icsatmes}) requires care, both because of differences in the assumed mean molecular mass and because Saturn's mean stratospheric temperature is strongly affected by seasonally varying insolation \\citep{BezardGautier1985icsatclimate}.  Heating by inertia-gravity waves might also be important, at least in some regions and/or seasons \\citep{CoorayEtal1998icnpsatoc, FrenchGierasch1974jasocwav, YoungEtal1997scijupgravwav}.  Detailed modeling of stratospheric temperatures requires taking account of non-LTE effects as well \\citep{Appleby1990icch4gp}.  A clearer picture of zonal and seasonal variations in Saturn's stratospheric structure should emerge when Earth-based stellar occultations can be viewed in the context of data from the {\\em Cassini} orbiter."
    },
    "1005/1005.4485_arXiv.txt": {
        "abstract": "Resistive relativistic magnetohydrodynamic (RRMHD) simulations are applied to investigate the system evolution of relativistic magnetic reconnection. A time-split Harten--Lan--van Leer method is employed. Under a localized resistivity, the system exhibits a fast reconnection jet with an Alfv\\'{e}nic Lorentz factor inside a narrow Petschek-type exhaust. Various shock structures are resolved in and around the plasmoid such as the post-plasmoid vertical shocks and the ``diamond-chain'' structure due to multiple shock reflections. Under a uniform resistivity, Sweet--Parker-type reconnection slowly evolves. Under a current-dependent resistivity, plasmoids are repeatedly formed in an elongated current sheet. It is concluded that the resistivity model is of critical importance for RRMHD modeling of relativistic magnetic reconnection. ",
        "introduction": "Magnetic reconnection \\citep{sweet,parker,petschek} is the driver of explosive events in space, astrophysical, and laboratory plasmas. The reconnection process attracts growing attentions to explain flaring events \\citep{lyut06,gia09} and the magnetic annihilation \\citep{coro90,lyu01} in relativistic plasma environments. Basic properties of relativistic magnetic reconnection have been discussed by relativistic magnetohydrodynamic (RMHD) theories \\citep{bf94b,lyu05,ten10}. In particular, \\citet{lyu05}'s careful work brought significant insights.  In the last decade, modern simulation works have revealed many features of relativistic magnetic reconnection. From the viewpoint of kinetic physics, it is widely recognized that reconnection is an efficient particle accelerator \\citep{zeni01,zeni07,claus04}. It was further found that the system is crucially influenced by the guide-field magnetic topology \\citep{zeni05} and the radiative cooling effects \\citep{claus09}. In a fluid scale, the reconnection system has been explored by a resistive RMHD (RRMHD) \\citep{naoyuki06} and relativistic two-fluid \\citep{zeni09a,zeni09b} models. \\citet{zeni09a} demonstrated and extensively analyzed a quasi-steady Petschek-type reconnection.  There is a strong demand for further development of RRMHD reconnection work. Unlike the kinetic and two-fluid models, the RRMHD model is free from kinetic scales such as the skin depth and the gyro radius, and so it is highly desirable to study stellar-scale problems. In addition, plasmas are considered to be collisional in optically thick, radiation-dominated environments \\citep{uz06}. Since such plasmas can be approximated by a single RMHD fluid, the RRMHD code with non-FLD\\footnote{FLD is flux-limited diffusion} radiative transfer will be necessary for future modeling of radiative relativistic reconnection.  However, no other RRMHD work has come out since \\citet{naoyuki06}, because the RRMHD equations turned out to be numerically unstable. Attempts to better solve the RRMHD system have been undertaken. \\citet{kom07} pointed out that the non-ideal electric field behaves as stiff relaxation to the ideal RMHD condition. He split the equations into two parts: the stiff terms were analytically solved as an exponential decay, and the rest part was solved by a two-step Harten--Lan--van Leer (HLL) method. \\citet{pal09} employed a hybrid scheme of an implicit Runge-Kutta method for the electric field and an explicit one for the other variables. \\citet{dumbser09} developed a higher-order scheme in unstructured volumes.  In this Letter, we present RRMHD simulations of relativistic magnetic reconnection, based on recent advances in numerical schemes \\citep{kom07}. ",
        "conclusions": "Our results corroborate the first RRMHD work by \\citet{naoyuki06} on many aspects such as Alfv\\'{e}nic outflow, the compression ratio, and the narrower slow shock structure. Judging from their inflow speed ($v_{in}/c_{A,in}\\sim \\mathcal{R}$), they observed faster reconnection probably due to the stronger localized resistivity of $S_1 \\sim \\mathcal{O}(1)$. Our results also agree with the two-fluid work by \\citet{zeni09a} such as the typical outflow parameters of the reference runs and the Petschek angle relations. Some visible differences are attributed to the two-fluid effects. The current layers look much sharper in this work, because the fluid inertial effects tend to smooth structures in the two-fluid model. The plasmoid looks longer in the $x$-direction. The outflow jet immediately hits the downstream plasmas in the $x$-direction in the RRMHD model, while plasma out-of-plane motion softens the reflection in the two-fluid model. Concluding this and those works, (1) Alfv\\'{e}nic outflow, (2) increasing reconnection rate, and (3) narrower opening angle, can be regarded as common features of relativistic Petschek reconnection. Since a narrower outflow exhaust suppresses the energy throughput, (3) does not immediately fit (2). However, \\citet{zeni09b} pointed out that the relativistic enthalpy flux transports huge energy per a cross section even in a narrower exhaust.  Thanks to the stable RRMHD code, we successfully resolved shock structures in and around the plasmoid. To our knowledge, the post-plasmoid vertical shocks and the diamond-chain structure are new discoveries. The diamond chain develops when the plasmoid edge speed (a sizable fraction of $c_{A,in}$) exceeds the sound speed in the neutral Harris sheet. We predict that it also develops in nonrelativistic MHD simulations once the above condition is met. Since multiple shocks are confined in a narrow region, the plasmoid edge can be a potential site of shock acceleration of particles.  From the numerical viewpoint, relativistic reconnection in the high-$\\sigma_{\\varepsilon}$ regimes is challenging for an explicit-type RRMHD code. Since the analytic part of the scheme contains the time constant of $-S\\gamma$ (e.g., Equation \\ref{eq:anom2}a), employing Alfv\\'{e}nic outflow speed, a code requires $ \\Delta t \\lesssim S_0^{-1} (1+\\sigma_{\\varepsilon,in})^{-1/2}$. As shown in Section 3, the outflow is faster near the plasmoid neck, and so the restriction becomes even severer. Thus, high time (and appropriate spatial) resolutions are necessary in high-$\\sigma_{\\varepsilon}$ and moderate-$S$ regimes of our interest. In addition, a plasma solution is not always stable even when the total momentum and energy are well handled. Let us extract the plasma momentum from Equation \\ref{eq:mom}, $\\partial_{t} \\vec{m} + \\div ( p\\vec{I} + w \\vec{u}\\vec{u} ) = \\vec{j} \\times \\vec{B}$. Focusing on the source term which is usually problematic, we see that the plasma stability condition, ${|\\Delta \\vec{m}|}/{w} \\sim {| \\vec{j} \\times \\vec{B} | \\Delta t }/{w} \\lesssim \\mathcal{O}(0.1),$ will be critical in higher-$\\sigma_{\\varepsilon}$ regimes, unless Ohm's law strongly limits the current. These issues need to be improved by an implicit approach \\citep{pal09}.  From the physics viewpoint, those restrictions obviously come from an immature form of Ohm's law. Note that the RRMHD model contains the displacement current $\\partial_t\\vec{E}$ in Amp\\`{e}re's law (Equation \\ref{eq:E}). Since it operates in a short timescale of the plasma frequency, the nonrelativistic MHD code usually drops it. However, the RRMHD system needs it for the energy and momentum conservation. On the other hand, the RRMHD uses a simple Ohm's law with the scalar resistivity (Equation \\ref{eq:ohm}). This is a time-stationary form \\citep{bf93} and there is no general consensus on extended forms. Combining the two equations with different timescales, the RRMHD system inevitably becomes problematic, especially in its impulsive phase. Thus, a further development of Ohm's law (e.g., \\citet{ged96}) or the reconstruction of the entire RRMHD system (e.g., \\citet{koide09}) is highly desirable.  In order to study large-scale problems which contain reconnection in arbitrary locations, we can no longer use the localized resistivity at a fixed location. We do need good physical or phenomenological anomalous resistivity models. Our attempt is a first step toward this direction, and the current-dependent resistivity model did exhibit different system evolution. This clearly tell us that the modeling of the effective resistivity, a long-standing problem of the entire reconnection research, is very important for RRMHD modeling of relativistic magnetic reconnection as well.  In summary, we explored an RRMHD model of relativistic magnetic reconnection. We corroborated earlier works and further studied fine plasmoid structures. Three different resistivity models are examined and they showed different system evolutions. It is crucially important to model an effective resistivity for RRMHD modeling of magnetic reconnection."
    },
    "1005/1005.4399_arXiv.txt": {
        "abstract": "% We study analytically the development of gravitational instability in an expanding shell having finite thickness. We consider three models for the radial density profile of the shell: (i) an analytic uniform-density model, (ii) a semi-analytic model obtained by numerical solution of the hydrostatic equilibrium equation, and (iii) a 3D hydrodynamic simulation. We show that all three profiles are in close agreement, and this allows us to use the first model to describe fragments in the radial direction of the shell. We then use non-linear equations describing the time-evolution of a uniform oblate spheroid to derive the growth rates of shell fragments having different sizes. This yields a dispersion relation which depends on the shell thickness, and hence on the pressure confining the shell. We compare this dispersion relation with the dispersion relation obtained using the standard thin-shell analysis, and show that, if the confining pressure is low, only large fragments are unstable. On the other hand, if the confining pressure is high, fragments smaller than predicted by the thin-shell analysis become unstable. Finally, we compare the new dispersion relation with the results of 3D hydrodynamic simulations, and show that the two are in good agreement. ",
        "introduction": "%  Shells and bubbles are common morphological features in the interstellar media of galaxies, and have been observed at many different wavelengths. For instance, \\citet{2002ApJ...578..176M} and \\citet{2005A&A...437..101E} have identified $\\sim 1000$ expanding shells in HI surveys of the Milky Way. HI shells have also been found in the LMC \\citep{1998ApJ...503..674K}, the SMC \\citep{2005MNRAS.360.1171H}, and other nearby galaxies. In the infrared, \\citet{2006ApJ...649..759C,2007ApJ...670..428C} have assembled a catalogue of $\\sim 600$ shells found in the GLIMPSE survey of our Galaxy, using the SPITZER space telescope. Shells have been detected in the Wisconsin H$\\alpha$ mapper data \\citep[WHAM;][]{2003ApJS..149..405H}. Shells have also been studied in millimetre molecular line emission, radio continuum and X-rays.  It is generally believed that a substantial fraction of these shells is formed by feedback from massive stars. In addition, it has been suggested by \\citet{1977ApJ...214..725E} that expanding shells can trigger star formation by the {\\it collect-and-collapse}  mechanism. In this mechanism, a massive star (or a group of stars) injects energy into the interstellar medium in the form of stellar winds and ionising radiation, and creates a bubble of gas with high temperature. The bubble expands due to its high internal pressure, and sweeps up the ambient medium into a dense shell. The shell cools down, becomes gravitationally unstable, fragments and forms a new generation of stars. If the new generation includes massive stars, the whole process may repeat itself, leading to sequential self-propagating star formation. Various modifications to the collect-and-collapse mechanism are reviewed by \\citet{1998ASPC..148..150E}.  The collect-and-collapse mechanism has been tested observationally in a series of papers by \\citet{2003A&A...408L..25D, 2005A&A...433..565D, 2006A&A...458..191D, 2008A&A...482..585D, 2009A&A...496..177D} and \\citet{2006A&A...446..171Z}. Using data at various wavelengths, they identify and study several objects in which star formation appears to have been triggered at the borders of HII regions. \\citet{2008ApJ...681.1341W,2009ApJ...694..546W} analyse six shells from the GLIMPSE survey, and determine the properties of young stellar objects concentrated in these shells.  The gravitational instability of an expanding infinitesimally thin shell has been studied theoretically by several authors. \\citet{1983ApJ...274..152V} derives a dispersion relation (perturbation growth rate as a function of wavenumber) for an infinitesimally thin shell, using decomposition into spherical harmonics. A very similar dispersion relation is obtained by \\citet{1994ApJ...427..384E} using perturbation analysis of linearised two-dimensional hydrodynamic equations. \\citet{1994MNRAS.268..291W} use the equation of motion for a two-dimensional fragment forming on the surface of a shell, and find expressions for the size of the first and most unstable fragment, and for the time at which the instability starts. All these analyses use a very similar physical model (an infinitesimally thin shell which expands into a uniform medium), and although the mathematical description in each of them is very different, the results regarding fragment sizes and growth rates are very similar.  We have tested these analytic predictions for the growth rate of gravitational instability, based on the thin-shell approximation, in \\citet[][hereafter Paper~I]{2009MNRAS.398.1537D}. We use two very different three-dimensional hydrodynamic codes (an Eulerian AMR code and a Lagrangian SPH code) to simulate the evolution and subsequent fragmentation of expanding self-gravitating shells. Our setup differs slightly from the one studied by Vishniac, Elmegreen and Whitworth, in that our shell does not accrete mass from the ambient medium. This modification requires only minor changes to the linear theory of the thin shell gravitational instability, which we describe in \\S\\ref{sec:thin_shell}.  We make this modification because an accreting shell is prone to the Vishniac dynamical instability \\citep{1983ApJ...274..152V}, which would grow quickly and make an analysis of gravitational instability impossible. On the other hand, we want to have the option to keep the shell thin, in order to test the thin-shell approximation. Therefore, we fill the shell interior and exterior with a hot rarefied gas and the pressure of this gas confines the shell. However, because of its low density, the ram pressure of the accreted gas, which is crucial for the development of the Vishniac instability, is negligible. We also keep the pressure of the rarefied gas constant throughout the evolution, so that the shell is effectively in free fall, and development of the Rayleigh-Taylor instability is suppressed.  In Paper~I, we find  excellent agreement between the two numerical codes. However, for simulations with low confining pressure, in which the shell becomes thick, the simulations do not agree well with the predictions of the thin-shell approximation. The agreement is better for simulations in which the confining pressure is chosen so that the shell thickness stays approximately constant during its evolution.  In this paper (Paper~II) we study the gravitational instability of a thick shell analytically, in order to understand how a thick shell fragments, and why the fragmentation differs from the predictions of the thin-shell approximation. We show that the confining pressure, which defines the shell thickness, is an important factor in determining the range of unstable wavelengths. Therefore, we call the new model {\\em pressure assisted gravitational instability} (PAGI).  The outline of the paper is as follows. In \\S\\ref{sec:radprof} we derive a semi-analytic description of the equilibrium radial profile of the shell, and compare it with our 3D numerical simulations, and with the simple uniform-density model of the shell that we adopt in the subsequent analysis. Section \\S\\ref{sec:gi} deals in detail with modelling gravitational instability in a shell. We briefly derive the dispersion relation for a non-acreting shell, using the thin-shell approximation; we give a description of a fragment forming in a shell with non-zero thickness, and compare it with the description used in the thin-shell approximation; finally we derive a new dispersion relation for a thick pressure-confined shell. In \\S\\ref{sec:num}, we compare the new dispersion relation with growth rates measured in our 3D numerical simulations, in \\S\\ref{discussion} we discuss the astrophysical consequences of our results, and in \\S\\ref{conclusions} we summarise our conclusions. ",
        "conclusions": "%  We have studied an isothermal ballistic shell, confined from both sides by a hot highly rarefied gas having non-zero pressure, $P_{_{\\rm EXT}}$. We have solved the equation of hydrostatic equilibrium in a frame of reference moving with the shell and shown that the resulting shell density profiles are in very good agreement with the profiles obtained from three-dimensional hydrodynamic simulations. We have also shown that the shell thickness measured at half of its peak density agrees well with the thickness of a simple uniform-density shell model.  This allows us to model a fragment forming in the shell, due to gravitational instability, as a uniform-density oblate spheroid. We use this model to derive a dispersion relation for pressure assisted gravitational instability (PAGI). This dispersion relation (perturbation growth rate as a function of initial fragment wavenumber) depends on the pressure of the external medium: the higher $P_{_{\\rm EXT}}$, the larger the maximum unstable wavenumber (i.e. the smaller the smallest unstable fragment), and the faster the growth rate for an unstable fragment. If $P_{_{\\rm EXT}}\\!=\\!0$, the highest unstable wavenumber is $0.6$ times smaller than predicted by the standard thin-shell analysis using the same parameters; if $P_{_{\\rm EXT}}\\!=\\!\\infty$, the highest unstable wavenumber is $2.2$ times larger. The PAGI dispersion relation gives approximately the same range of unstable wavelengths as the thin-shell dispersion relation if $P_{_{\\rm EXT}}\\simeq P_{_{\\rm CRIT}}$, where $P_{_{\\rm CRIT}}$ is the critical external pressure for which self-gravity and external pressure contribute equally to confinement of the shell.  Finally, we have demonstrated that the predictions of the PAGI dispersion relation agree rather well with the results of three-dimensional hydrodynamic simulations. In particular, (i) the PAGI dispersion relation predicts a maximum unstable wavenumber very similar to the simulations (but systematically $\\sim 25\\%$ higher); (ii) modulo this offset, the increase in $l_{_{\\rm PAGI}}^{^{\\rm MAX}}$ with $P_{_{\\rm EXT}}$ predicted by the PAGI dispersion relation is exactly mirrored by the simulations."
    },
    "1005/1005.1269_arXiv.txt": {
        "abstract": "We present two \\textit{Suzaku} observations of the Galactic center microquasar \\mq~separated by approximately 700 days. The source was observed on both occasions after a transition to the spectrally hard state. Significant emission from \\mq~ is detected out to an energy of 300 keV, with no spectral break or turnover evident in the data. We tentatively measure a lower limit to the cut-off energy of $\\sim$ 380 keV.  The spectra are found to be consistent with a Comptonized corona on both occasions, where the high energy emission is consistent with a hard power-law ($\\rm \\Gamma \\sim 1.8$) with a significant contribution from an accretion disc with a temperature of $\\sim$ 0.4 keV at soft X-ray energies.  The measured value for the inner radius of the accretion disc is found to be inconsistent with the picture whereby the disc is truncated at large radii in the low-hard state and instead favours a radius close to the ISCO ($\\rm R_{in} \\sim 10 - 20~R_g$). ",
        "introduction": "Since the discovery of apparently super-luminal jets from the X-ray binary GRS 1915+105 \\citep{b12}, the Galactic microquasars have assumed a position of critical importance in our efforts to understand accretion physics and relativistic jet production \\citep{b40}.  \\mq~ was discovered by the \\textit{Einstein} satellite \\citep{b1}. Subsequent observations revealed \\mq~to be the dominant source of hard X-rays ($\\rm > 20~keV$) in the direction of the Galactic center \\citep{b2,b3}, where the source is located approximately 50 arcmin from Sgr A$\\rm ^*$.  The microquasar nature of \\mq~was discovered upon the observation of a double sided radio jet consistent with the X-ray position \\citep{b4}.  Further VLA observations showed this radio source to be highly variable \\citep{b5}.  Since its discovery \\mq~has been observed on numerous occasion at X-ray wavelengths. The column density towards this source is high given the proximity to the Galactic center and has been measured by \\textit{Chandra} to be $\\rm \\sim 1 \\times 10^{23}~cm^{-2}$ \\citep{b10}. Here, the spectrum was found to be consistent with a power-law ($\\rm \\Gamma \\sim 1.4$). \\textit{INTEGRAL} low-hard state observations have detected \\mq~up to energies of $\\sim$ 600 keV, where the spectrum is found to be consistent with a power-law ($\\rm \\Gamma \\sim 1.6$) up to 200 keV, with an additional component required at higher energies \\citep{b8}, see also \\citet{b9} for earlier \\textit{INTEGRAL/RXTE} observations. In the high-soft state, the hard X-ray flux decreases significantly, dropping below the \\textit{INTEGRAL} detection limit at energies $\\rm >$ 50 keV \\citep{b8}. \\citet{b11} reported on 5 years of \\textit{RXTE} monitoring, where they discovered a modulation with a period 12.73$\\rm \\pm$0.05 days which is attributed to the orbital period of the binary, in addition to a possible super-orbital modulation with a period of $\\sim$ 600 days.  As the extinction at optical wavelengths is prohibitive ($\\rm A_v \\sim 50$), counterpart searches are required to take place in the infrared. While a number of candidate counterparts have been identified \\citep{b6,b7}, no variability has been observed from these, rendering them unlikely to be the actual counterpart. The current upper limit for the K$\\rm_s$-band magnitude of the counterpart is $\\geq$ 19.9 at the 95\\% confidence level. This is equivalent to a secondary spectral type of O or B if on the main sequence or K if a giant, at an assumed distance of 8.5 kpc.  In this paper, we describe observations undertaken with the \\textit{Suzaku} X-ray observatory, while \\mq~was in the low-hard state. In \\S2, we describe the observations and extraction of source spectra. We proceed to analyze the data in \\S3. In \\S4, these results are compared to observations of other microquasars in the hard state, and finally our conclusions are presented in \\S5. ",
        "conclusions": "We present \\textit{Suzaku} observations of the Galactic center microquasar \\mq~ in two separate epochs taken after the system had transitioned into the low-hard state. The system is observed to be in the low-hard state at the time of our observations with an X-ray luminosity of $\\sim$ 1\\% Eddington.  The spectra in each epoch are similar, being described by a model consisting of a soft thermal accretion disc component (T $\\sim$ 0.4 keV) and the broadband emission ($\\rm > 10 keV$) is found to be characterized by an unbroken power-law to at least 300 keV.  Consistent with growing evidence from observations of numerous systems in the low-hard state (e.g. GX 339-4, SWIFT J1753.5-0127, XTE J1817-330, XTE J1118+480), we also find evidence that the accretion disc in \\mq~is not truncated at large radii in the low-hard state and instead resides close to the ISCO ($\\rm R_{in} \\sim 20~R_g$).\\\\"
    },
    "1005/1005.1433_arXiv.txt": {
        "abstract": "{ Herschel PACS and SPIRE images have been obtained over a 30\\arcmin$\\times$30\\arcmin\\ area around the well-known carbon star CW Leo (IRC +10 216). An extended structure is found in an incomplete arc of $\\sim$22\\arcmin\\ diameter, which is cospatial with the termination shock due to interaction with the interstellar medium (ISM) as defined by Sahai \\& Chronopoulos from ultraviolet GALEX images. Fluxes are derived in the 70, 160, 250, 350, and 550 $\\mu$m bands in the region where the interaction with the ISM takes place, and this can be fitted with a modified black body with a temperature of 25$\\pm$3~K. Using the published proper motion and radial velocity for the star, we derive a heliocentric space motion of 25.1 \\ks. Using the PACS and SPIRE data and the analytical formula of the bow shock structure, we infer a de-projected standoff distance of the bow shock of $R_{0} = (8.0 \\pm 0.3) \\times 10^{17}$ cm. We also derive a relative velocity of the star with respect to the ISM of $(106.6 \\pm 8.7)/\\sqrt{n_{\\rm ISM}}$ \\ks, where n$_{\\rm ISM}$ is the number density of the local ISM. } ",
        "introduction": "Ever since the discovery paper by Becklin et al. (1969), the object IRC +10 216 (= AFGL 1381 = CW Leo) has spurred much interest. We now know that it is a carbon star in an advanced stage of stellar evolution on the asymptotic giant branch (AGB), pulsating and surrounded by an optically thick dust shell and large molecular circumstellar envelope (CSE). One aspect of study has been to constrain the properties of the CSE by answering questions such as what is the mass-loss rate and how has it changed with time, what kind of chemistry takes place, and what is the geometry and structure of the CSE?   The deep optical images taken by Mauron \\& Huggins (1999, 2000), Mauron et al. (2003), and Le\\~{a}o et al. (2006) show that the dusty envelope is not smooth but consists of a series of arcs or incomplete shells. The average angular separation between the dust arcs suggests a timescale for the change in mass-loss rate of the order of 200--800 yr.  The lack of kinematic information on the dust arcs precludes any firm conclusion about the true three-dimensional structure of the arcs or shells.  From large-scale mapping at a relatively low angular resolution of the CO J=1-0 emission, Fong et al. (2003) discovered a series of large molecular arcs or shells at radii of $\\sim$100\\arcsec\\ in the outer envelope. They attribute these multiple shells as ``being the reverberations of a single Thermal Pulse erupting over 6000 yr ago.'' The timescale inferred from the spacing between these arcs is about 200--1000 yr.  In the present paper, we discuss our new results on the outer shell of CW Leo from observations with the {\\em Herschel Space Observatory} (Pilbratt et al. 2010) and their connection to the results from {\\em GALaxy Evolution Explorer Space Observatory} (GALEX) by Sahai \\& Chronopoulos (2010, hereafter SC). For the present analysis, we adopted a distance of $d=$ 135\\,pc and a mass-loss rate \\mdot= 2.2$\\cdot 10^{-5}$ \\msolyr\\ (Groenewegen et al. 1998), a gas expansion velocity of V$_{\\rm exp}$ = 15.4 \\ks, a radial velocity V$_{\\rm LSR}= -25.5$ \\ks\\ (Groenewegen et al. 2002), corresponding to V$_{\\rm helio}= -18.6$\\,\\ks, and a proper motion (pm) $\\mu_{\\rm \\alpha} \\cos \\delta = +26 \\pm$6, $\\mu_{\\rm \\delta} =+4 \\pm$6 mas/yr (Menten et al. 2006). ",
        "conclusions": "\\label{discussion}  \\subsection{Bow shock, thermal pulse, or both ?}  The most widely accepted explanation of large detached shells is mass--loss variation (e.g., Olofsson et al. 1990, Zijlstra et al. 1992). AGB stars experience thermal pulses (TPs) during which intense mass loss ejections occur. A star can undergo several TPs, which would lead to the formation of concentric spherical detached shells (see Kerschbaum et al. 2010).  Another explanation of detached envelopes is the interaction between the AGB wind and the interstellar medium (ISM) (e.g. Young et al. 1993, Martin et al. 2007, Ueta et al. 2006).  In this scenario, the AGB wind is slowed down as it sweeps up material from the ISM. The piled up material forms a density enhancement that continue to expand due to the internal pressure. Shocks can occur if the relative velocity of the AGB wind with respect to the ISM is large enough. The thermal emission of the dust in the density enhancement at the shock interface between the stellar wind and the the ISM can be detected in the far-IR (Ueta et al. 2006, 2009). While a detached shell produced by TPs would be spherical, in the case of wind-ISM interaction the shape of the shell will depend on the space motion of the star through the ISM. A wind--ISM shell can look spherical as seen for R Cas (Ueta et al. 2009) if most of the space motion of the star relative to the ISM is in the radial direction. For stars with a high space motion (relative to the ISM), the shape of the bow shock will be more parabolic with the apex of the parabola in the direction of the star's motion relative to the ISM.  In the case of CW\\,Leo, the shape of the extended emission and the position of the star suggest that the stellar wind has driven a shock into the ISM. The far-IR emission is probably caused by the thermal emission of the piled-up dust at the shock interface.  One may think that the observed UV emission is produced by dust scattering of the interstellar light. However, the brightness ratio of FUV to NUV is $\\sim$10 (see Table~\\ref{TableF}), which is much higher than expected in that case (of the order of $\\sim$2.4; SC). The only other AGB star with UV data probing a wind--ISM interaction is Mira (Martin et al. 2007).  Martin et al. suggest that collisional excitation of cool H$_2$ by hot electrons from the post-shock gas is responsible for the UV emission. The faint NUV emission can be explained by the H$_2$ emitting in the FUV band. SC suggest that the same mechanism may also be the dominant contributor to the FUV ring emission in CW Leo. The dust and the FUV emission have a similar spatial scale. For the planetary nebula NGC 6720, we note that van Hoof et al. (2010) find on the basis of PACS and SPIRE data  that dust and H$_2$ are co-spatial and argue that H$_2$ has been formed on grain surfaces.  \\begin{figure} \\begin{center} \\vspace{0cm} \\hspace{0cm} \\resizebox{9cm}{!}{ \\includegraphics[angle=90]{figures/14658fig2.ps}} \\caption{Intensity profiles as a function of the offset along the minor axis of the ellipse increasing from west to east. The profile for each wavelength was normalised to the intensity at 558$\\arcsec$ and shifted up for clarity. From bottom to top with normalisation factor in mJy/arcsec$^2$ and the vertical shift: FUV (0.065,0.0), 160~$\\mu$m (0.063,0.8), 250~$\\mu$m (0.023,1.1), 350~$\\mu$m (0.008,1.4), 500~$\\mu$m (0.004,1.8). The vertical dotted line indicates the position of the dust emission peak from the center of the ellipse.} \\label{Ipro} \\vspace{-4mm} \\end{center} \\end{figure}   The structure seen in our data may also be the result of both mechanisms i.e., mass--loss variation and wind--ISM interaction. By estimating the interpulse period for CW Leo, we can check wether the dust shell we see is from an earlier TP that is now interacting with the ISM.  Guelin et al. (1995) constrained the initial mass of CW Leo to be 3 \\less\\ $M$ \\less\\ 5 \\msol\\, based on the isotopic ratios of $^{24,25,26}$Mg. From the initial--final mass relation from Salaris et al. (2009), this implies a likely final mass (and essentially the current core mass) of 0.7--0.9 \\msol. From the core mass interpulse period relation of Wagenhuber \\& Groenewegen (1998) for solar metallicity stars, this implies an interpulse period of  6~000--33~000 years.   Scaling the mass loss and the wind velocity in SC with our assumed values, the flow timescale in the unshocked and shocked winds is 19~900 and 56~000 yr (very long because of the very low velocity of 1.2~\\ks\\ in this region, see SC for details), for a total lower limit to the duration of the mass-loss phase of about 75~000 years. If TPs were to modulate the mass loss, as hypothesised for the origin for the detached shells discussed earlier, one could expect at least one other TP to have occurred during the time it took the envelope to expand this far. No obvious density enhancement is evident in the unshocked wind in the PACS and SPIRE images, which suggests that the interpulse period is at least 19~000 years. However, the dynamical and interpulse timescales are compatible, so it is possuble that the bow shock and dust emission are not only the result of a steady outflow interacting with the ISM, but might include the effect of an enhanced wind of short duration due to a recent TP (or pulses).  \\begin{table}[!h] \\caption{Derived fluxes for the extended emission} \\begin{tabular}{cccccc} \\hline \\hline $\\lambda$  &  Flux & $\\lambda$ &  Flux & $\\lambda$ &  Flux    \\\\ ($\\mu$m)   &  (Jy) & ($\\mu$m)   &  (Jy) & ($\\mu$m)   &  (Jy) \\\\ \\hline 0.15 & $(20\\pm0.2)\\times10^{-4}$  &  70  & 3.51$\\pm$0.54 &  250 & 4.31$\\pm$0.66 \\\\ 0.23 & $ (2\\pm0.2)\\times10^{-4}$ & 160  & 7.70$\\pm$1.17 &  350 & 1.70$\\pm$0.27 \\\\ &                          &      &               &  550 & 0.73$\\pm$0.14 \\\\ \\hline \\end{tabular} \\label{TableF} \\vspace{-7mm} \\end{table}  \\subsection{Space motion of the star and the ISM flow velocity}  From the adopted radial velocity, distance, and proper motion, the heliocentric space velocity of the star is derived to be about 25.1 \\ks\\ at a heliocentric inclination angle of the space motion vector of 47.8 degrees (measured from the plane of the sky, away from us) and a PA of 81.3\\degr\\ for the proper-motion vector in the plane of the sky.  In the PACS/SPIRE wavebands, the observed far-IR surface brightness is expected to be proportional to the column density of the dusty material in the shell because the optical depth of the shell is much lower than unity.  Because the bow shock interface is a parabolic surface arbitrarily oriented in space, the column density tends to reach its highest value where the bow shock cone intersects with the plane of the sky including the central star.  Thus, the apparent shape of the bow shock is the conic section of the bow shock.  Therefore, given the analytical formula of the bow shock structure (Wilkin 1996), we can fit the observed surface brightness of the bow shock to the conic section of the bow shock cone to derive the heliocentric orientation of the bow shock (e.g. Ueta et al. 2008; 2009).  From this fitting, we determined that the apex of the bow shock cone is oriented at $61\\fdg9 \\pm 0\\fdg3$ (this is degenerate, in a sense that it could be pointing away from us or towards us) with respect to the plane of the sky into the PA of $88^{\\circ}$ with the deprojected stand-off distance of $(8.0 \\pm 0.3) \\times 10^{17}$ cm.  Using these numbers in the ram pressure balance equation $V_{\\star} = \\sqrt{\\frac{\\dot{M} \\; V_{\\rm exp}}{4 \\pi \\; \\mu_{\\rm H} \\; m_{\\rm H} \\; n_{\\rm ISM} \\; R_0^2} }$ for a mean nucleus number per hydrogen nucleus of $\\mu_{\\rm H}= 1.4$, the relative velocity of the star with respect to the ISM is $V_{\\star} = (106.6 \\pm 8.7)/\\sqrt{n_{\\rm ISM}}$ \\ks, where $n_{\\rm ISM}$ is the number density of the ISM local to CW Leo in cm$^{-3}$.   By following the scheme of Johnson \\& Soderblom (1987, also see Ueta et al. 2008), the heliocentric space velocity components of the star can be converted to the heliocentric Galactic space velocity components [U, V, W] of [21.6$\\pm$3.9, 12.6$\\pm$3.5, -1.8$\\pm$3.3] \\ks\\ , and also to the LSR Galactic space velocity components of [30.6$\\pm$3.9, 24.6$\\pm$3.5, 5.2$\\pm$3.3] \\ks. The heliocentric ISM flow velocity is 117.6 \\ks\\ if the bow cone is facing us (i.e.~the apex pointing toward) or 82.6 \\ks\\ if the bow cone is facing away from us.   \\begin{figure}[t] \\begin{center} \\vspace{0cm} \\hspace{0cm} \\resizebox{9cm}{!}{ \\includegraphics[angle=90]{figures/14658fig3.ps}} \\caption{Modified black body ($T$= 25$\\pm$3~K and $\\beta$= 1.6$\\pm$0.4) fit (dashed curve) to the derived fluxes (symbols) for the extended emission. } \\label{dustT} \\vspace{-4mm} \\end{center} \\end{figure}   Wareing et al. successfully modelled bow shocks for AGB stars using 3-D hydrodynamical models (R\\,Hya and Mira in Wareing et al. 2006, 2007a). Their models were able to reproduce in great detail all the components of a bow shock (astropause, astrotail, vortices) and constrain the space velocity of the star, the ISM density, and the mass--loss rate. The models were built by varying the velocity of the star relative to the ISM, the ISM density and the mass loss rate and the effect of those physical parameters on the shape of the bow shock can be seen in Wareing et al. (2007b). From the Wareing models, we can see that for ISM densities $\\le$2~cm$^{-3}$ (which implies $V_{\\star} \\ge$75 \\ks) and after an evolution of 50~000~yr into the AGB phase, the bow shock does not resemble our data. The models have a more flattened bow shock with the star closer to the apex of the shock. This is for a mass--loss rate between 5$\\times 10^{-7}$~\\msolyr\\ and 10$^{-6}$~\\msolyr, which is an order of magnitude lower than the value assumed for CW\\,Leo. For a higher mass--loss rate, we would expect an even more important departure from sphericity with a significant tail of ejecta.  This suggests that the density of the local ISM to CW\\,Leo is probably higher than 2~cm$^{-3}$ implying an upper limit of 75 \\ks to $V_{\\star}$, the stellar velocity relative to the ISM."
    },
    "1005/1005.2606_arXiv.txt": {
        "abstract": "Sky coverage is one of the most important pieces of information about astronomical observations. We discuss possible representations, and present algorithms to create and manipulate shapes consisting of generalized spherical polygons with arbitrary complexity and size on the celestial sphere. This shape specification integrates well with our Hierarchical Triangular Mesh indexing toolbox, whose performance and capabilities are enhanced by the advanced features presented here. Our portable implementation of the relevant spherical geometry routines comes with wrapper functions for database queries, which are currently being used within several scientific catalog archives including the Sloan Digital Sky Survey, the Galaxy Evolution Explorer and the Hubble Legacy Archive projects as well as the Footprint Service of the Virtual Observatory. ",
        "introduction": "\\label{sec:intro}  Astronomers have to keep accurate records of where their observations are located on the sky. Beyond the direction and angular size of the field of view, we have detailed information available about the precise sky coverage derived from the exact shape of our detectors. This coverage is invaluable for most statistical studies, e.g., luminosity functions or especially analyses of spatial clustering.  The de facto standard of the Flexible Image Transport System \\citep[FITS;][]{fits} has reserved keywords for the World Coordinate System \\citep[WCS;][]{wcs} specification, and parameters that specify the transformation from image pixels to sky coordinates (and reverse), are present in the header of most FITS files. While the WCS is perfectly adequate for individual exposures, multiple observations are difficult to describe in a single system. Every field has potentially a separate coordinate system, hence moving from field to field is convoluted, and makes it cumbersome to answer even simple questions, e.g., whether two separate fields overlap.  The footprint of a large-area survey might be complicated but the small-scale irregularities are even more problematic. Not only the depth of a survey changes as a function of the position on the sky, but parts of the observations are often censored for various reasons, such as bright stars blocking the view, satellite trails, artifacts from reflections. If we would like to represent all these on the sky, we need a scalable solution that works for shapes of arbitrary complexity and size from the subpixel level to the entire sky. We need tools to create and manipulate these descriptions right there where the data are and utilize the information efficiently.  Geographic Information Systems (GIS) were designed with a similar goal in mind. There are however subtle differences, that are large enough that GIS systems are not quite applicable to astronomy directly. The modern mapping systems do not extend much beyond the basic features of projected maps but provide very efficient tools for finding nearby places and shortest routes, etc. Even the most complicated GIS shapes are limited to spherical polygons whose sides are great circles (or straight lines in the projection.) In astronomy, the approximations with such polygons would be unacceptably inaccurate or prohibitively redundant, hence there is need for an extended set of features to represent the geometries of the observations and surveys.  Some of the concepts discussed in the paper are not new and have been introduced and studied previously in different fields. Books by \\citet{samet89,samet90} detail the quad-trees for spatial searches that \\citet{barrett94} advocated using for astronomy data. \\cite{fekete90a,fekete90b,short95} developed an icosahedron-based methodology for Earth sciences, and \\citet{kunszt00, kunszt01} introduced the convex description and the Hierarchical Triangular Mesh, a triangulation that was also used by \\citet{lee98} with a different numbering scheme. A similar representation of shapes is also found in \\citet{mangle1} and \\citet{mangle2}. \\citet{goodchild91,goodchild92} and \\citet{song00} introduced the Discrete Global Grid for GIS systems, an equal area variant of the same triangulation idea. The integration to relational databases is discussed in \\citet{gray04}.   The goal this paper is to provide the astronomy community with a complete and consistent view of the current and much improved methodology built on a new fully functional spherical geometry framework, and describe the implementations and interfaces used in several astronomy archives and tools today including the Sloan Digital Sky Survey \\citep{cas}, the Hubble Legacy Archive \\citep{greene07}, the Galaxy Evolution Explorer and the NVO Footprint Service \\citep{adass_footprint}. In Section~\\ref{sec:shapes} we describe how to specify spherical shapes and manipulate them. Section~\\ref{sec:geom} deals with the spherical geometry of the region representations and Section~\\ref{sec:htm} discusses an efficient search method based on the Hierarchical Triangular Mesh. In Section~\\ref{sec:lib} we provide details of our software solution including the implementation of the database routines and their usage. Section~\\ref{sec:sum} summarizes the main results, along with current and future applications in astronomy. ",
        "conclusions": ""
    },
    "1005/1005.5026_arXiv.txt": {
        "abstract": "The mass of the Galactic dark matter halo is under vivid discussion. A recent study by Xue et al. (2008, ApJ, 684, 1143) revised the Galactic halo mass downward by a factor of $\\sim$2 relative to previous work, based on the line-of-sight velocity distribution of $\\sim$2400 blue horizontal-branch (BHB) halo stars. The observations were interpreted in a statistical approach using cosmological galaxy formation simulations, as only four of the 6D phase-space coordinates were determined. Here we concentrate on a close investigation of the stars with highest negative radial velocity from that sample. For one star, SDSSJ153935.67+023909.8 (J1539+0239 for short), we succeed in measuring a significant proper motion, i.e. full phase-space information is obtained. We confirm the star to be a Population~II BHB star from an independent quantitative analysis of the {\\em SDSS} spectrum -- providing the first NLTE study of any halo BHB star -- and reconstruct its 3D trajectory in the Galactic potential. J1539+0239 turns out as the fastest halo star known to date, with a Galactic rest-frame velocity of 694$^{+300}_{-221}$\\,\\kms\\ (full uncertainty range from Monte Carlo error propagation) at its current position. The extreme kinematics of the star allows a significant lower limit to be put on the halo mass in order to keep it bound, of $M_{\\rm halo}\\ge1.7_{-1.1}^{+2.3}\\times10^{12}$\\,\\msun. We conclude that the Xue et al. results tend to underestimate the true halo mass as their most likely mass value is consistent with our analysis only at a level of 4\\%. However, our result confirms other studies that make use of the full phase-space information.\\\\ ",
        "introduction": "\\label{sec:intro} Knowledge of the properties of dark matter halos is an important issue for our understanding of galaxy formation and evolution, and for unveiling the nature of dark matter. The halo of the Milky Way therefore is of highest interest, as it allows unique observational constraints to be obtained for testing theoretical models \\citep[e.g.][]{1996ApJ...462..563N}. Several observational campaigns -- e.g. the Sloan Digital Sky Survey \\citep[\\sdss,][]{2000AJ....120.1579Y}, the RAdial Velocity Experiment \\citep[\\emph{RAVE},][]{2006AJ....132.1645S}, the Sloan Extension for Galactic Understanding and Exploration \\citep[\\emph{SEGUE},][]{2009AJ....137.4377Y}  -- provide the tracers for studying halo properties, like the total mass of the halo and its extent.  Several studies in the past decade determined the halo mass from ever increasing samples of halo stars, globular clusters and/or satellite galaxies. However, it is in fact only a few objects at highest velocity that are affecting largely the mass estimates \\citep{2003A&A...397..899S,2007MNRAS.379..755S}. While larger halo masses of about $2\\times10^{12}$\\,\\msun\\ were favoured earlier \\citep{1999MNRAS.310..645W,2003A&A...397..899S}, lower masses of about half this value were derived more recently \\citep{2005MNRAS.364..433B,2007MNRAS.379..755S,2008ApJ...684.1143X}. The precise value determines, e.g. among others, whether satellite galaxies like the Magellanic Clouds \\citep[e.g.][]{2006ApJ...652.1213K,2009AJ....137.4339C} or hyper-velocity star\\footnote{Hyper-velocity stars move at such high velocity that they may be gravitationally unbound to the Galaxy. Their supposed place of origin is the Galactic center, where they may have been accelerated by gravitational interactions with the super-massive black hole \\citep{hills88}.} candidates \\citep{2009ApJ...691L..63A} are on bound orbits, or not.  Most of the previous studies had to rely substantially on the {\\em distributions} of radial velocities in their samples to derive their conclusions as full space motions for halo objects are unavailable in many cases at present. In such cases only four coordinates (i.e. two position values, distance and radial velocity, RV) of the 6D phase space are determined and the missing data (the proper motion components) are handled in a statistical approach. E.g., radial velocities for more than 10,000 blue halo stars from the \\sdss were measured in the yet most extensive study by \\citet{2008ApJ...684.1143X}. The sample was composed of blue horizontal-branch, blue straggler, and main-sequence stars with effective temperatures roughly between 7,000 and 10,000\\,K according to their colours. Here, we focus on the fastest stars of the Xue et al. sample in terms of {\\em negative} line-of-sight velocity, indicating a {\\em bound} orbit.  For one of them we were able to measure a significant proper motion, which allowed a detailed 3D kinematic investigation when combined with a quantitative spectroscopic analysis that facilitated the determination of the star's distance. \\object{SDSSJ153935.67+023909.8} (J1539+0239 for short) is an inbound\\footnote{An extragalactic origin of the star as an unbound low-mass hyper-velocity star from another galaxy is imaginable but unlikely. The local volume is devoid of galaxies in a wide range (several 10\\,$\\deg$) around the infall direction of J1539+0239 at present \\citep[for the Local Group see][and later discoveries]{1999A&ARv...9..273V}. If ejected early after its formation, the star could have travelled $\\sim$8\\,Mpc during its lifetime, such that a final answer on this may be only obtained when full space motions of the nearby galaxies are known.} Population~II horizontal branch star with a Galactic rest-frame (GRF) velocity of $\\sim$700\\,\\kms\\ at its current position, making it the fastest halo object known. This allows a significant lower limit for the total halo mass of the Galaxy to be set. The present work provides a glimpse to the detailed kinematic investigations feasible once the European Space Agency's Gaia satellite mission \\citep[e.g.][]{2005ESASP.576.....T} becomes operational, at much higher precision. ",
        "conclusions": "\\label{sec:conclusion}  We reported the quantitative spectral analysis of a high-velocity star from the sample of faint blue halo stars of \\citet{2008ApJ...684.1143X}. J1539+0239 was confirmed to be a Population~II blue horizontal branch star with a low metallicity of $[$Fe/H$]=-2.0$ and the characteristic enhancement of $\\alpha$-elements. Hereby, we performed a NLTE analysis of a halo BHB star for the first time. While the majority of the weak lines were confirmed to be formed close to LTE conditions, many of the stronger metal lines -- which are important diagnostics at the spectral resolution achieved within the {\\em SDSS} -- show pronounced NLTE strengthening, with the differences between the derived LTE and NLTE abundances amounting to 0.1\\,dex to 0.5\\,dex typically. In addition to information on the chemical composition, the radial velocity, proper motion and spectroscopic distance were derived and a detailed kinematical analysis was performed.  Carrying out kinematical numerical experiments using the Galactic potential of \\citet{1991RMxAA..22..255A} in order to obtain an orbit of J1539+0239 gravitationally bound to the Milky Way, we found that the mass of the dark halo has to be at least $M_{\\rm halo}\\sim1.7_{-1.1}^{+2.3}\\times10^{12}$\\,\\msun (absolute uncertainties from extrema in MC error propagation) . This mass limit is in good agreement with several previous studies \\citep{1999MNRAS.310..645W,2003A&A...397..899S,2009ApJ...691L..63A}. However, the significantly lower most likely mass value of \\cite{2008ApJ...684.1143X} is consistent with our analysis only at a level of 4\\%, i.e. it likely underestimates the Galactic dark halo mass.  We conclude, that if the kinematics of a halo star is extraordinary enough, and the errors within the analysis are small, even {\\em one} star alone can provide a significant lower limit to the dark matter halo mass, and to the total mass of the Milky Way (halo\\,+\\,bulge\\,+\\,disk), here $M_{\\rm total}\\ge1.8_{-1.1}^{+2.3}\\times10^{12}$\\,\\msun. The determining factor is that full kinematic information is available, as it will become routine in the era of the Gaia space mission, at much higher precision."
    },
    "1005/1005.5356_arXiv.txt": {
        "abstract": "{} {We determine the physical properties (spin state and shape) of asteroid (21) Lutetia, target of the ESA Rosetta mission, to help in preparing for observations during the flyby on 2010 July 10 by predicting the orientation of Lutetia as seen from Rosetta.} {We use our novel KOALA inversion algorithm to determine the physical properties of asteroids from a combination of optical lightcurves, disk-resolved images, and stellar occultations, although the latter are not available for (21) Lutetia.} {We find the spin axis of (21) Lutetia to lie within 5\\degr~of ($\\lambda = 52$\\degr, $\\beta = -6$\\degr) in Ecliptic J2000 reference frame (equatorial $\\alpha = 52$\\degr, $\\delta = +12$\\degr), and determine an improved sidereal period of 8.168\\,270 $\\pm$ 0.000\\,001 h. This pole solution implies the southern hemisphere of Lutetia will be in ``\\textsl{seasonal}'' shadow at the time of the flyby. The apparent cross-section of Lutetia is triangular as seen ``\\textsl{pole-on}'' and more rectangular as seen ``\\textsl{equator-on}''. The best-fit model suggests the presence of several concavities. The largest of these is close to the north pole and may be associated with large impacts. } {} ",
        "introduction": "\\indent The origin and evolution of the Solar System and its implications for early planetesimal formation are key questions in planetary science. Unlike terrestrial planets, which have experienced significant mineralogical evolution, through endogenic activity, since their accretion, small Solar System bodies have remained essentially unaltered. Thus, a considerable amount of information regarding the primordial planetary processes that occurred during and immediately after the accretion of the early planetesimals is still present among this population. Consequently, studying asteroids is of prime importance in understanding the planetary formation processes \\citep{2002-AsteroidsIII-1-Bottke} and, first and foremost, requires a reliable knowledge of their physical properties (size, shape, spin, mass, density, internal structure, etc.) in addition to their compositions and dynamics. Statistical analyses of these parameters for a wide range of asteroids can provide relevant information about inter-relationships and formation scenarios. \\\\ \\indent In this respect, our observing program with adaptive optics, allowing diffraction-limited observations from the ground with 10 m-class telescopes, has now broken the barrier which separated asteroids from real planetary worlds \\citep[\\textsl{e.g.,}][]{2007-Icarus-191-Conrad, 2008-AA-478-Carry, 2009-Icarus-202-Drummond, 2010-Icarus-205-Carry-a, 2010-AA--Drummond}. Their shapes, topography, sizes, spins, surface features, albedos, and color variations can now be directly observed from the ground. This opens these objects to geological, rather than astronomical-only, study. While such surface detail is only possible for the largest asteroids, our main focus is on determining accurately the size, shape, and pole. Among them, we have observed (21) Lutetia, an asteroid that will be observed in-situ by the ESA Rosetta mission.\\\\ \\indent The Rosetta Mission will encounter its principal target, the comet 67P/Churyumov-Gerasimenko, in 2014. However, its interplanetary journey was designed to allow close encounters with two main-belt asteroids: (2867) \\v{S}teins and (21) Lutetia. The small asteroid (2867) \\v{S}teins was visited on 2008 September 5 at a minimum distance of about 800 km \\citep{2009-DPS-40-Schulz} and (21) Lutetia will be encountered on 2010 July 10. Knowing the geometry of the flyby (\\textsl{e.g.}, visible hemisphere, sub-spacecraft coordinates as function of time, and distance) before the encounter is crucial to optimize the observation sequence and schedule the on-board operations. The size of Lutetia \\citep[estimated at $\\sim$100 km, see][]{ 2002-AJ-123-Tedesco, PDSSBN-Iras, 2006-AA-447-Mueller} allows its apparent disk to be spatially resolved from Earth. Our goal is therefore to improve knowledge of its physical properties to prepare for the spacecraft flyby. \\\\ \\indent Lutetia, the Latin name for the city of Paris, is a main-belt asteroid (semi-major axis 2.44 AU) that has been studied extensively from the ground \\citep[see][for a review, primarily of recent observations]{2007-SSRv-128-Barucci}. Numerous studies have estimated indirectly its spin \\citep[by lightcurve, \\textsl{e.g.,}][]{ 1987-KFNT-3-Lupishko-a, 1992-AA-95-Dotto, 2003-Icarus-164-Torppa}. Size and albedo were reasonably well determined in the 1970s by \\citet{1977-Icarus-31-Morrison} using thermal radiometry (108--109 km), and by \\citet{1976-AJ-81-Zellner} using polarimetry (110 km). Five somewhat scattered IRAS scans \\citep[\\textsl{e.g.,}][]{2002-AJ-123-Tedesco, PDSSBN-Iras} yielded a higher albedo and smaller size than the dedicated observations in the 1970s. \\citet{2006-AA-447-Mueller} derived results from new radiometry that are roughly compatible with the earlier results or with the IRAS results, depending on which thermal model is used. \\citet{2008-AA-479-Carvano} later derived a lower albedo from ground-based observations, seemingly incompatible with previous works. Radar data analyzed by \\citet{1999-Icarus-140-Magri, 2007-Icarus-186-Magri} yielded an effective diameter for Lutetia of 116 km; reinterpretation of those data and new radar observations \\citep{2008-Icarus-195-Shepard} suggest an effective diameter of 100 $\\pm$ 11 km and an associated visual albedo of 0.20. Recent HST observations of Lutetia \\citep{2009-arXiv-Weaver} indicate a visual albedo of about 16\\%, a result based partly on the size/shape/pole determinations from our work in the present paper and \\citet{2010-AA--Drummond}.\\\\ \\indent Lutetia has been extensively studied using spectroscopy in the visible, near- and mid-infrared and its albedo measured by polarimetry and thermal radiometry \\citep{1975-ApJ-195-McCord, 1975-Icarus-25-Chapman, 1976-AJ-81-Zellner, 1978-Icarus-35-Bowell, 1995-Icarus-117-Rivkin, 1999-Icarus-140-Magri, 2000-Icarus-145-Rivkin, 2004-AA-425-Lazzarin, 2005-AA-430-Barucci, 2006-AA-454-Birlan, 2007-AA-470-Nedelcu, 2008-AA-477-Barucci, 2008-Icarus-195-Shepard, 2009-Icarus-202-DeMeo, 2009-Icarus-202-Vernazza, 2009-AA-498-Lazzarin, lazzarin-2010, 2010-AA-513-Perna, 2010-AA-Belskaya}. We present a discussion on Lutetia's taxonomy and composition in a companion paper \\citep{2010-AA--Drummond}.\\\\ \\indent Thermal infrared observations used to determine the size and albedo of  Lutetia were initially inconsistent, with discrepancies in diameters and visible albedos reported \\citep[\\textsl{e.g.,}][]{% 1976-AJ-81-Zellner, 1989-Icarus-78-Lupishko, 1996-PSS-44-Belskaya, 2002-AJ-123-Tedesco, 2006-AA-447-Mueller, 2008-AA-479-Carvano}. \\citet{2006-AA-447-Mueller} and \\citet{2008-AA-479-Carvano}, however, interpreted these variations as an indication of surface heterogeneity, inferring that the terrain roughness of Lutetia increased toward northern latitudes\\footnote{our use of ``northern hemisphere'' refers to the hemisphere in the direction of the positive pole as defined by the right-hand rule from IAU recommendations \\citep{2007-CeMDA-98-Seidelmann}}, that the crater distribution is different over the northern/southern hemispheres, and includes a possibility of one or several large craters in Lutetia's northern hemisphere. Indeed, the \\textsl{convex} shape model derived from the inversion of 32 optical lightcurves \\citep{2003-Icarus-164-Torppa} displays a flat top near the north pole of Lutetia. \\citet{2002-AA-383-Kaasalainen} have shown that large flat regions in these convex models could be a site of concavities. The southern hemisphere is not expected to be free from craters however, as \\citet{2010-AA-513-Perna} detected a slight variation of the visible spectral slope, possibly due to the presence of large craters or albedo spots in the southern hemisphere.\\\\ \\indent In this paper, we present simultaneous analysis of adaptive-optics images obtained at the W. M. Keck and the European Southern Observatory (ESO) Very Large Telescope (VLT) observatories, together with lightcurves, and we determine the shape and spin state of Lutetia. In section~\\ref{sec: obs}, we present the observations, in section~\\ref{sec: shape} the shape of Lutetia, and finally, we describe the geometry of the upcoming Rosetta flyby in section~\\ref{sec: flyby}. ",
        "conclusions": "} \\indent We have reported disk-resolved imaging observations of (21) Lutetia obtained with the W. M. Keck and Very Large Telescope observatories in 2007, 2008, and 2009. We have derived the shape and spin of (21) Lutetia using the Knitted Occultation, Adaptive-optics, and Lightcurve Analysis (KOALA) method, which is based on combining these AO images with optical lightcurves gathered from over four decades.\\\\ \\indent The shape of (21) Lutetia is well described by a Camembert wedge, and our shape model suggests the presence of several concavities near its north pole and around its equator. The spin axis of Lutetia is tilted with respect to its orbital plane, much like Uranus, implying strong seasonal effects on its surface. At the time of the Rosetta flyby, (21) Lutetia's northern hemisphere will be illuminated while the southern hemisphere will be in long-term darkness, hindering the size determination from Rosetta.\\\\ \\indent The next opportunity to observe Lutetia's shortest dimension, impacting its volume determination, will occur in 2011 July, one year after Rosetta flyby, when the sub-Earth point will be close to its equator (SEP$_\\beta$ of +31\\degr). During this time, observations using large telescopes equipped with adaptive-optics will allow refinement of Lutetia's short dimension and thus improve the volume determination. This ground-based support will be essential to take advantage of the high-precision mass determination provided by the spacecraft deflection observed during the flyby."
    },
    "1005/1005.0868_arXiv.txt": {
        "abstract": "In Brans-Dicke theory a non-linear self interaction of a scalar field $\\phi$ allows a possibility of realizing the late-time cosmic acceleration, while recovering the General Relativistic behavior at early cosmological epochs. We extend this to more general modified gravitational theories in which a de Sitter solution for dark energy exists without using a field potential. We derive a condition for the stability of the de Sitter point and study the background cosmological dynamics of such theories. We also restrict the allowed region of model parameters from the demand for the avoidance of ghosts and instabilities. A peculiar evolution of the field propagation speed allows us to distinguish those theories from the $\\Lambda$CDM model. ",
        "introduction": "Scalar-tensor gravitational theories have been widely studied as an alternative to General Relativity (GR). In these theories a scalar-field degree of freedom $\\phi$ is coupled to the Ricci scalar $R$ through a coupling of the form $F(\\phi)R$ \\cite{FujiiMaeda}. For example, such a coupling arises in low energy effective string theory as a result of the dilaton coupling with gravitons \\cite{GasVene}. The discovery of dark energy in 1998 \\cite{SNIa} has also stimulated the study for the modifications of gravity on large distances (see Refs.~\\cite{darkreview} for recent reviews).  The well-known example of scalar-tensor theories is Brans-Dicke (BD) theory \\cite{BDtheory} in which $F(\\phi)$ is proportional to $\\phi$ with a non-canonical kinetic term $(-\\omega_{\\rm BD}/\\phi)(\\nabla \\phi)^2$, where $\\omega_{\\rm BD}$ is the so-called BD parameter. In original BD theory without a field potential, the BD parameter is constrained to be $\\omega_{\\rm BD}>40000$ from local gravity tests in the solar system \\cite{Will}. This comes from the fact that the coupling between the massless field $\\phi$ and non-relativistic matter needs to be suppressed to avoid the propagation of the fifth force. Under this bound, the deviation from GR (which corresponds to the limit $\\omega_{\\rm BD} \\to \\infty$) is too small to be detected in current cosmological observations.  There are two ways to recover the General Relativistic behavior in high density regimes relevant to solar system experiments. One is the so-called chameleon mechanism \\cite{chame} in which the presence of the field potential $V(\\phi)$ allows a possibility of having a density-dependent effective mass of the field. Provided that the effective mass is sufficiently large in the regions of high density, a spherically symmetric body can have a thin-shell around its surface so that the effective coupling between $\\phi$ and matter is suppressed outside the body. This mechanism works not only for BD theory \\cite{chame,TUMTY} but also for metric $f(R)$ gravity \\cite{fRchame}, because the latter can be regarded as a special case of BD theory with $\\omega_{\\rm BD}=0$ \\cite{ohanlon,Chiba}. For viable dark energy models based on BD theory and metric $f(R)$ gravity the early cosmological evolution mimics that of GR, but the deviation from GR becomes important at late times (i.e. low density regimes) \\cite{TUMTY,fRviable,fRviable2}.  Another way for the recovery of GR in the regions of high density is to introduce non-linear self interactions of a scalar field, e.g., $\\xi (\\phi) \\square \\phi (\\partial^{\\mu} \\phi \\partial_{\\mu} \\phi)$, where $\\xi$ is a function in terms of $\\phi$. There have been attempts to restrict the form of the Lagrangian by imposing the ``Galilean'' symmetry $\\partial_{\\mu} \\phi \\to \\partial_{\\mu} \\phi +b_{\\mu}$ \\cite{Nicolis,Deffayet,Chow,Rham,Gannouji}. The self interaction of the form $\\square \\phi (\\partial^{\\mu} \\phi \\partial_{\\mu} \\phi)$ respects the Galilean symmetry in the Minkowski background. In the Dvali-Gabadadze-Porrati (DGP) braneworld model \\cite{DGP} the field self interaction $\\square \\phi (\\partial^{\\mu} \\phi \\partial_{\\mu} \\phi)$ has been employed to recover the General Relativistic behavior for the length scale smaller than the so-called Vainshtein radius \\cite{DGPse1,DGPse2} (see also Refs.~\\cite{Babichev}). While this ``self screening'' mechanism (called the Vainshtein mechanism \\cite{Vain}) can be at work for consistency with solar system experiments, the DGP model is unfortunately plagued by a ghost problem \\cite{DGPse2,DGPghost} as well as incompatibility with a number of observational constraints \\cite{DGPobser}.  In BD theory with the field self interaction term $\\xi (\\phi) \\square \\phi (\\partial^{\\mu} \\phi \\partial_{\\mu} \\phi)$, there exist de Sitter (dS) solutions for $\\xi(\\phi) \\propto \\phi^{-2}$ and $\\omega_{\\rm BD}<-4/3$ \\cite{Silva}\\footnote{For constant $\\xi (\\phi)$ there exist no consistent dS solutions, but the presence of other field self interaction terms that respect Galilean symmetry in the Minkowski background allows a possibility for giving rise to dS solutions \\cite{Gannouji}.}. Moreover the General Relativistic behavior is recovered in early cosmological epochs during which the field is nearly frozen. The solutions finally approach the dS attractor at which $\\dot{\\phi}/\\phi$ is constant. As long as $\\dot{\\phi}/\\phi$ is positive, one can avoid the appearance of ghosts and instabilities associated with the field propagation speed. Moreover this model gives rise to several interesting observational signatures such as the modified growth of matter perturbations and anti-correlations in the cross-correlation of large scale structure and the integrated Sachs-Wolfe effect in cosmic microwave background anisotropies \\cite{Silva,Kobayashi1}.  In this paper we consider the general action (\\ref{action}) below without the field potential and derive the functional forms of $F(\\phi)$, $B(\\phi)$, and $\\xi (\\phi)$ from the requirement of obtaining dS solutions with $\\dot{\\phi}/\\phi=$\\,constant. The three functions are restricted to be of the power-law forms in terms of $\\phi$, which include BD theory with $\\xi (\\phi) \\propto \\phi^{-2}$ as a special case. We discuss the cosmological viability of such theories by analyzing the stability of fixed points for the background field equations. We also study conditions for the avoidance of ghosts and instabilities to find the region of viable model parameters. ",
        "conclusions": "\\label{conclude}  We have constructed modified gravitational models of dark energy starting from the general action (\\ref{action}) without a field potential. The presence of a non-linear self interaction $\\xi(\\phi) \\square \\phi (\\partial^{\\mu} \\phi \\partial_{\\mu} \\phi)$ allows us to recover the General Relativistic behavior in the regions of high density. The functions $F(\\phi)$, $B(\\phi)$, and $\\xi(\\phi)$ are restricted in the power-law forms given in Eq.~(\\ref{funchoice}) from the demand of obtaining dS solutions at late times. The theory with $n=2$ corresponds to Brans-Dicke theory with the self-interaction $\\xi(\\phi) \\propto \\phi^{-2}$, which was recently studied in literature \\cite{Silva}. If $\\xi(\\phi)=0$, the theories with the functions $F(\\phi)=M_{\\rm pl}^2 (\\phi/M_{\\rm pl})^{3-n}$ and $B(\\phi)=\\omega ( \\phi/M_{\\rm pl})^{1-n}$ are equivalent to Brans-Dicke theory by introducing a new field $\\chi \\equiv F(\\phi)$. However, except for $n=2$, the presence of the field self interaction term does not allow us to express theories with the functions (\\ref{funchoice}) as Brans Dicke theory plus the term $\\xi(\\chi) \\square \\chi (\\partial^{\\mu} \\chi \\partial_{\\mu} \\chi)$ by such a field redefinition.  We have derived the dynamical equations (\\ref{auto1})-(\\ref{auto3}) to discuss the background cosmological evolution for the theories with the functions (\\ref{funchoice}). The evolution of the dimensionless variables $x$ and $y$ during radiation and matter eras can be analytically estimated, which is consistent with the results obtained by numerical integrations. The presence of these epochs demands the condition $n \\le 3$. The variables $x$ and $y$ at the dS point, denoted as $x_{\\rm dS}$ and $y_{\\rm dS}$ respectively, are determined by solving Eqs.~(\\ref{xdS}) and (\\ref{ysolu}) for given $n$ and $\\omega$. We also studied the stability of the dS point by considering the evolution of homogenous curvature perturbations and found that the stability condition is given by $(n-3)x_{\\rm dS}<3$.  The viable model parameter space can be restricted further by studying conditions for the avoidance of ghosts and instabilities. The no-ghost condition corresponds to $Q_s>0$, where $Q_s$ is defined in Eq.~(\\ref{Qdef}). Provided $F(\\phi)>0$, we require that $x>0$ to avoid ghosts during the cosmological evolution from the radiation era to the dS epoch. We note that the case $n=2$ is special because of the divergent behavior of the quantity $Q_s$ at the dS point. Interestingly the stability of the dS point is automatically satisfied for $n \\le 3$ and $x_{\\rm dS}>0$.  The instability of perturbations can be avoided for $c_s^2>0$, where the propagation speed squared $c_s$ is defined in Eq.~(\\ref{cses}). At the dS point $c_s^2$ reduces to Eq.~(\\ref{eq:csdS}), which is positive for $n \\ge 2$ under the conditions $n \\le 3$ and $x_{\\rm dS}>0$. Hence the viable parameter region of $n$ is constrained to be $2 \\le n \\le 3$. For the existence of the dS point the parameter $\\omega$ is restricted in the range $\\omega<-n(n-3)^2$ from Eq.~(\\ref{xdS}). When $n \\neq 3$ the evolution of $c_s^2$ during radiation and matter dominated epochs can estimated as Eqs.~(\\ref{csra}) and (\\ref{csma}) respectively, which remain subluminal. During the transition from the matter era to the dS epoch the propagation speed can be superluminal, depending on the values of $\\omega$ and $n$. The avoidance of the temporal superluminal propagation gives the lower bound on $\\omega$ for each value of $n$. In Fig.~\\ref{omegafig} we have plotted the viable model parameter space and the region of the subluminal propagation in the $(n, \\omega)$ plane. When $n=3$ it is difficult to avoid the appearance of the superluminal mode during the matter era because of the specific evolution of the variables $x$ and $y$.  It will be of interest to study the evolution of matter density perturbations and gravitational potentials to confront our theories with the observations of large scale structure, cosmic microwave background, and weak lensing. We leave the detailed analysis of cosmological perturbations for future work."
    },
    "1005/1005.1163_arXiv.txt": {
        "abstract": "We present the first  multi-fluid analysis of a dense neutron star core with a deconfined colour-flavour-locked superconducting quark component. Accounting only for the condensate and (finite temperature) phonons, we  make progress by taking over results for superfluid $^4$He. The resultant two-fluid model accounts for a number of additional viscosity coefficients (compared to the Navier-Stokes equations) and we show how they enter the dissipation analysis for an oscillating star. We provide simple estimates for the gravitational-wave driven r-mode instability, demonstrating that the various phonon processes that we consider are not effective damping agents. Even though the results are likely of little direct astrophysical importance (since we  consider an overly simplistic stellar model)  our analysis represents significant technical progress, laying the foundation for more detailed numerical studies and  preparing the ground for the inclusion of additional aspects (in particular associated with kaons) of the problem. ",
        "introduction": "The state of matter at extreme densities continues to be an issue of vigorous investigation. The problem is complicated, not only from the theoretical point of view, but also by the fact that laboratory experiments are restricted. For example, while colliders like RHIC at Brookhaven, GSI in Darmstadt and the LHC at CERN probe hot quark-gluon plasmas they will never be able to explore the cold, extreme high density, region of the QCD phase diagram. In order to test our understanding of the relevant physics we need to turn to astrophysics, and the dynamics of compact stars. In fact, ``neutron stars'' represent unique laboratories of such extreme physics. With core densities  reaching about one order of magnitude beyond nuclear saturation, they are likely to contain exotic states of matter like hyperon phases with net strangeness and/or deconfined quarks. It is well-established that these states of matter should exhibit superfluidity/superconductivity at the relevant temperatures (neutron stars are born with temperatures $\\sim 10^{12}$~K and rapidly cool below $\\sim 10^9$~K). Moreover, observed radio pulsar glitches provide strong evidence for the presence of a partially decoupled superfluid component in these systems. The modelling of the dynamics of these, potentially \\underline{very} complex, objects presents a serious challenge.  A key aspect of the problem concerns the fact that a superfluid system has additional dynamical degrees of freedom. This is well-known from experiments on laboratory systems like $^4$He, which exhibit a second sound associated with thermal waves \\cite{khalatnikov,putterman}. Analogous ``superfluid'' modes have been studied in detail for superfluid neutron-proton-electron mixtures relevant for the outer core of a neutron star \\cite{epstein,mendell,ac01}. For simplicity, these studies have almost exclusively ignored thermal effects (the work in \\cite{gusakov} is a notable exception). While this is a useful first approximation it is clear that the zero temperature assumption must be relaxed in a realistic model. Basically, due to the density dependence of the various superfluid pairing gaps (see \\cite{NPA} for a guide to the literature), there will \\underline{always} be regions in a neutron star where thermal effects are important (in the vicinity of the critical density at which the phase transition occurs). Understanding the nature of these transition regions, and their effect on various aspects of neutron star dynamics, is one of the main challenges for  research in this area.  In this paper, we describe a first attempt at modelling thermal dynamics in a superfluid neutron star. We focus on stars with a colour-flavour-locked (CFL) superconducting quark core \\cite{CFL} at finite temperatures. This is an interesting problem for several reasons. First of all, the simplest possible model for this system considers a quark condensate coupled to a gas of phonons. This problem is analogous to $^4$He at low temperatures, and hence we can bring our recent dissipative two-fluid model \\cite{helium} to bear on it (more or less directly). The lessons we learn from this exercise should inform the development of finite temperature models for the superfluids in the outer core and the neutron star crust. Secondly, even though there have been discussions of the dynamics of the different CFL phases, in particular in the context of the gravitational-wave driven r-mode instability (see \\cite{jaikumar} for references), the superfluid aspects have (as far as we are aware) not previously been accounted for. Our discussion begins to address the relevance of the additional degrees of freedom in these systems, and  provides some insight into the nature of the different ``fluids'' involved.  As an application with immediate astrophysical relevance, we will work out the inertial r-modes and estimate the relevant viscous damping rates for the ``simplest'' model of CFL matter. We consider the ``cool'' regime where the temperature is significantly below all the quasiparticle energy gaps. In this regime, dissipation may mainly occur due to phonon interactions. One reason for considering this model is that there are results in the literature for both bulk- and shear viscosity \\cite{manuel1,manuel2} as well as the mutual friction associated with superfluid vortices \\cite{mana1}. Of course, the simple ``condensate plus phonon'' model that we  consider is not the whole story. It should apply at asymptotically high densities, but may not be the true ground state at lower densities. The discussion in \\cite{alford1,alford2,alford3} adds extra dimensions to the problem by considering the bulk viscosity due to kaons (allowing for flavour-changing processes), which will be present at higher temperatures. At first sight, this mechanism will only be relevant for very hot stars ($T>1$~MeV~$\\sim 10^{10}$~K) since it assumes that there is a thermal population of kaons. Below the critical energy where the kaons appear the contribution to the bulk viscosity is exponentially suppressed and may not be that important. However, in more recent work on the so-called CFL-K$^0$ phase \\cite{alford2}, it is argued that the main low temperature mechanism involves condensed kaons. This is important since the kaon condensate will remain present as $T \\to 0$. The upshot of this is that the problem requires a ``multi-fluid'' analysis at all temperatures. The simplest model would  have three components;  the quark condensate, the kaon condensate and finite temperature excitations (phonons and thermal kaons). This problem is  more involved than the case that we focus on here. Nevertheless, our results provide an essential starting point for investigations into the dynamical role of the kaons. Most importantly, by providing the ``hydrodynamics'' view of the problem we illustrate the input needed to study various dynamical scenarios. This should stimulate further discussion between experts on different aspects of this multi-faceted problem, as required to make progress in the future. ",
        "conclusions": "We have presented the first true multi-fluid analysis of a dense neutron star core with a deconfined, colour-flavour-locked superconducting, quark core. By focussing on a cool system, and accounting only for the condensate and (finite temperature) phonons, we  made progress by taking over much of the formalism from the analogous problem for superfluid $^4$He, the archetypal two-fluid laboratory system. The additional fluid degree of freedom, in the present case represented by the phonon gas, leads to the system not being well represented by the Navier-Stokes equations. In particular, a complete model requires a number of additional viscosity coefficients.  Without an actual calculation it is not easy to establish whether the multi-fluid aspects are relevant or not. For example, in the case of the gravitational-wave driven instability of the f-modes it is known that the superfluid degree of freedom is very important, since the vortex mutual friction may completely suppress that instability below a critical temperature (see \\cite{lm1,fmode}). It is known that, because of the different nature of the associated velocity field, the mutual friction does not affect the r-mode instability in the same drastic fashion \\cite{rmode,lm}. These  examples provide clear evidence that different problems need to be considered on a case-by-case basis.  We have provided a detailed dissipative formulation for a system comprising a quark (CFL) condensate and phonons. The model builds on recent improvements in our understanding of the analogous problems of superfluid $^4$He \\cite{helium} and causal heat conductivity \\cite{heat}. A key ingredient is the massless entropy component that represents the phonon excitations. We have discussed how the superfluid constraint of irrotationality reduces the number of required viscosity coefficients to four (one shear and three bulk), and provided a translation between  results in the literature \\cite{manuel1,manuel2} and the coefficients in our formalism. We also emphasised that many more dissipative channels may come into play in a rotating system where superfluid vortices are present \\cite{monster,helium}. In particular, we showed how the vortex mediated mutual friction is accounted for in the model, and translated the available results for the associated coefficients \\cite{mana1}.  In order to be able to make relevant estimates for the r-mode instability, we developed a simple two-component equation of state based on the MIT bag model at zero temperatures  with an additional phonon gas representing the thermal component. This example highlights the additional information that is needed in a multi-fluid analysis, in particular, regarding the entrainment coupling. Future work needs to provide a consistent equation of state, including \\underline{all} the key aspects. The model equation of state completed the formulation of the problem and we could, in principle, have carried out a numerical study of the r-modes. We opted not to do this, instead introducing a sequence of simplifying assumptions, because we felt that it would be useful to start by working out some less precise estimates for the relevant dissipation timescales. A numerical analysis of the problem should, of course, be encouraged. The problem has a number of interesting aspects, and may shed light on how one should deal with finite temperature superfluid neutron stars in general.  This work was motivated by a desire to understand the different phases of CFL matter from a hydrodynamics point of view. It is, obviously, an interesting problem and the notion that observations of gravitational waves from relativistic stars may help shed light on the extreme QCD phase diagram is exciting. The r-mode instability has been discussed in this context for some time, see for example \\cite{manuel1,manuel2,mana1,alford1,alford2,madsen,strange}, but there has not been any previous discussion of the multi-fluid aspects of the problem. We hope that this work will stimulate a more detailed discussion between experts in the relevant areas.  The fact that the various phonon processes that we have accounted for can be shown to have little effect on the r-mode instability should not discourage future efforts. In fact, the result could probably have been anticipated. The real challenge will be to account for additional degrees of freedom, especially associated with the kaons (either thermal or in a condensate) \\cite{alford1,alford2}, and the modelling of ``hybrid'' stars with quark cores and various phase transitions. These problems have additional features that, while possibly understood in principle, have never been considered in practice. This makes the modelling more complex, but also  intriguing since the richer dynamics may  lead to surprises."
    },
    "1005/1005.3020_arXiv.txt": {
        "abstract": "Recent observations of excited CO emission lines from \\zsim 2 disc galaxies have shed light on the \\schmidt \\ relation at high-\\z \\ via observed \\sigmacojtwo \\ and \\sigmacojthree \\ relations.  Here, we describe a novel methodology for utilising these observations of high-excitation CO to derive the underlying Schmidt (\\schmidt) relationship.  To do this requires an understanding of the potential effects of differential CO excitation with SFR.  If the most heavily star-forming galaxies have a larger fraction of their gas in highly excited CO states than the lower SFR galaxies, then the observed molecular Kennicutt-Schmidt index, $\\alpha$, will be less than the underlying \\schmidt \\ index, $N$. Utilising a combination of SPH models of galaxy evolution and molecular line radiative transfer, we present the first calculations of CO excitation in \\zsim 2 disc galaxies with the aim of developing a mapping between various observed \\sigmaco \\ relationships and the underlying \\schmidt relation.  We find that even in relatively luminous \\zsim 2 discs, differential excitation does indeed exist, resulting in $\\alpha < N$ for highly excited CO lines. This means that an observed (e.g.) \\sigmacojthree \\ relation does not map linearly to a \\sigmahtwo \\ relation. We utilise our model results to provide a mapping from $\\alpha$ to $N$ for the range of Schmidt indices $N=1-2$.  By comparing to recent observational surveys, we find that the observed \\sigmacojtwo \\ and \\sigmacojthree \\ relations suggest that an underlying SFR $\\propto \\rho^{1.5}$ relation describes \\zsim 2 disc galaxies. ",
        "introduction": "\\label{section:introduction}  Since the original works by \\citet{sch59} and \\citet{ken98b} parameterising star formation rates (SFRs) in galaxies in terms of the scaling relation: \\begin{equation} \\label{equation:schmidt} SFR \\propto \\rho_{\\rm gas}^{\\rm N} \\end{equation} there have been considerable efforts by both the Galactic and extragalactic star formation communities to characterise the exponent $N$ \\citep[e.g. ][and references therein]{ken98a}. Constraining the Schmidt SFR relation in galaxies is desireable, both for understanding the physics of star formation on local scales, as well as for giving simulators recipes for modeling processes below typical numerical resolution scales.  Because the volume density is not observable, most observed forms of Equation~\\ref{equation:schmidt} have been in terms of the SFR and gas surface densities\\footnote{To remain consistent with typical literature nomenclature, we will refer to the volumetric form of Equation~\\ref{equation:schmidt} as the Schmidt relation, and the surface density form as the Kennicutt-Schmidt (KS) relation.  We will reserve the index, $N$, for the volumetric exponent, and $\\alpha$ for the surface density exponent.}.  As a result of recent pioneering high resolution millimetre-wave surveys, two trends have become apparent.  First, the SFR surface density in galaxies is well correlated with the {\\it molecular} gas surface density, though little relation exists between the SFR and HI atomic gas \\citep[e.g. ][]{big08,ler08,kru09}. Second, the {\\it global} integrated relationship between SFR and CO molecular gas surface densities in observations of local galaxies carries an exponent $N \\sim 1.5$ \\citep[e.g. ][]{san91,ken98b,gao04a,gao04b}.  A variety of theories have been proposed to explain the $N \\sim 1.5$ index inferred from observations \\citep[e.g. ][]{sil97,tan00,elm02a,kru05,kru09b}, most of which rely on stars forming on a time scale proportional to the dynamical time scale.  The recent advent of sensitive (sub)mm-wave interferometers has allowed, for the first time, observational constraints on the \\schmidt \\ relation in ``normal'' star-forming (e.g. not starbursting) galaxies at high redshift via the detection of CO rotational emission lines \\citep[][]{bou07,ion09,dad10a,dad10b,bot09,gen10,tac10}.  However, owing to fact that most mm-wave interferometers operate in the 1-3 mm wavelength regime, most detections of CO emission lines at high-redshift are of relatively excited lines (e.g. CO (J=3-2), as opposed to the ground state transition which is typically used as a tracer of \\htwo \\ molecular gas.  In principle, if the excitation of CO is relatively invariant in a given sample of galaxies, one can simply make an assumption regarding the CO (J=3-2) to (J=1-0) scaling ratio in high-\\z \\ galaxies, and derive an \\htwo \\ gas mass (or surface density) with the inferred CO (J=1-0) intensity.  However, if the excitation of CO in galaxies is not constant with increasing SFR, then the exponent in the \\sfrco \\ relation may not naturally translate to an exponent in the \\schmidt \\ relation.  In other words, $\\alpha$ does not always equal $N$ when probing high excitation CO lines.  To see this, consider a sample of galaxies which are forming stars at a rate according to SFR $\\propto \\rho^{1.5}$.  If the CO (J=1-0) line faithfully traces the \\htwo \\ gas mass, then one could expect an observed relationship $\\Sigma_{\\rm SFR} \\propto \\Sigma_{\\rm CO J=1-0}^{1.5}$ \\citep{kru07,nar08b}.  However, if the galaxies with the highest SFR have a higher fraction of their CO gas excited into the CO J=3 state than the lowest SFR galaxies, one will observe a {\\it flatter} relationship between SFR and L$_{\\rm CO 3-2}$ than index $\\alpha=1.5$.  The \\sigmacojthree \\ relation will only map linearly to the \\sigmahtwo \\ relation if the excitation of CO is invariant with SFR.  In practice, this can only happen if the CO gas is thermalised in the observed lines (if a given excited state is in LTE).  This effect of differential molecular excitation on observed molecular SFR scaling relations has been observed in the local Universe.  The SFR-CO (J=1-0) relationship has an index $\\alpha \\approx 1.5$, while the SFR-CO (J=3-2) relationship has a flatter index $\\alpha \\approx 0.9$ \\citep{san91,yao03,nar05,ion09,bay09}.  Similarly, while the SFR-HCN (J=1-0) relationship is linear in local galaxies, the SFR-HCN (J=3-2) index is decidedly sublinear \\citep[with index $\\alpha \\sim 0.7$; ][]{gao04a,gao04b,bus08,gra08,jun09}.  Observations of individual star-forming clumps (which are massive enough to host stellar clusters) have been inconclusive regarding whether these global trends extend to smaller scales \\citep[][]{wu05,wu10}.  In the absence of a more direct tracer of \\htwo \\ gas than high-excitation CO, the potential effects of differential molecular excitation with SFR need to be quantified in order to derive the underlying SFR relation in high-redshift galaxies.  The few multi-line constraints of excitation in high-\\z \\ galaxies hints that CO may be subthermally excited even in the most luminous \\zsim 2 systems \\citep{wei07,dan09,car10,har10}.  This indicates that applying a uniform mapping from (e.g.) CO (3-2) to CO (J=1-0) line intensities will indeed be problematic. In this arena, numerical models can offer guidance.  Our aim in this paper is to calculate the mapping of observed \\sigmaco \\ relations of excited lines (e.g. CO J=2-1 and CO J=3-2) to an underlying \\schmidt \\ relation controlling the star formation.  In \\citet{nar09,nar10b,nar10a} and \\citet{hay10}, we have developed a merger-driven model for the formation of high-redshift ULIRGs which shows reasonable correspondence with observed SEDs, CO emission properties and number counts (C. Hayward et al. in prep.).  Here, we utilise the (idealised) progenitor disc galaxies of these model mergers to represent the star-forming discs at high-\\z \\ typically observed in CO emission line surveys \\citep[e.g. ][]{tac10,dad10b}.  We combine these hydrodynamic simulations of disc galaxies with 3D non-LTE molecular line radiative transfer calculations in order to calculate the full statistical equilibrium excitation properties of the molecules.  These methods allow us to determine the differential excitation of CO of \\zsim 2 disc galaxies with respect to SFR, and derive a mapping of an observed molecular Kennicutt-Schmidt law (\\sigmaco) to a \\schmidt \\ relationship. ",
        "conclusions": "\\begin{table*} \\centering \\caption{This Table provides the mapping between observed Kennicutt-Schmidt molecular surface-density indices, $\\alpha$, and underlying volumetric Schmidt indices, $N$.  We provide the mapping for CO transitions J=1-0 through J=4-3.  The numbers contained here constitute the information necessary to re-create a plot like Figure~\\ref{figure:ks_molslope} for any of the four modeled CO transitions and will aid interpretation of future observations of varying molecular transitions.  The ``errors'' denote a 25\\% uncertainty level which encompasses variations in the final solution upon changing initial conditions in our simulation (primarily the equation of state and normalisation of the SFR relation).} \\begin{tabular}{@{}c||cccc@{}} \\hline Schmidt Index $N$ & SFR-CO (J=1-0)$^\\alpha$ & SFR-CO (J=2-1)$^\\alpha$ & SFR-CO (J=3-2)$^\\alpha$ & SFR-CO (J=4-3)$^\\alpha$ \\\\ \\hline   1  & 0.96$\\pm 0.24$   &  0.89$\\pm0.22$   &0.76$\\pm 0.19$   &0.49$\\pm 0.12$\\\\ 1.5& 1.47 $\\pm 0.37$  &  1.36$\\pm0.34$   &1.08$\\pm 0.27$   &0.77$\\pm 0.19$\\\\ 2&   2.13 $\\pm 0.53$   &  1.95$\\pm 0.48$  &1.45$\\pm 0.36$   &0.95$\\pm 0.24$\\\\   \\hline \\end{tabular} \\label{table:indices} \\end{table*}       The principle result of this study is that, due to subthermal excitation in high-lying CO lines in high-\\z \\ galaxies, observed \\sigmaco \\ relations do not necessarily map linearly to \\sigmahtwo \\ relations.  This is part of a broader theoretical framework first developed by \\citet{kru07} and \\citet{nar08b} which posits that the underlying density distribution of galaxies is crucial in mapping the underlying Schmidt relation to observed molecular Kennicutt-Schmidt relations.  In this paper, we have expanded upon these studies by exploring the relationship between Schmidt and Kennicutt-Schmidt indices in models which aim to serve as reasonable representations of the \\zsim 2 star-forming disc galaxies being uncovered in sensitive optical/NIR observations \\citep[e.g. ][]{dad05,for09}.  Indeed, the models studied here have been shown in previous publications to satisfy the star-forming {\\it BzK} colour-selection criteria \\citep{nar10b}, as well as serve as reasonable progenitors for luminous, merger-driven \\zsim 2 submillimetre and 24 \\micron-selected galaxies \\citep{hay10,nar10a}.  We have additionally explored the dependence of the observed molecular Kennicutt-Schmidt index on varying underlying Schmidt relations. Our models, combined with the general theoretical framework of \\citet{kru07} and \\citet{nar08b} suggest that observed \\zsim 2 discs are subject to differential CO excitation with respect to SFR, and that the observed \\sigmacojtwo and \\sigmacojthree \\ relations may map to an underlying Schmidt index of $N=1.5$ controlling the SFR.  An important verifying aspect to these models is that they are able to explain the multitude of observed SFR-\\lmol \\ relations in the local Universe.  For example, a linear relationship has been observed between SFR and HCN (J=1-0) in local galaxies, a trend which has often been interpreted as evidence that a linear {\\it dense gas} Schmidt relation controlled the SFR.  Our models suggest the linear SFR-HCN (J=1-0) relation is in reality a combined effect of an underlying \\schmidt \\ index of $N=1.5$ and differential excitation in HCN \\citep{kru07,nar08b}.  This view has been observationally confirmed by \\citet{bus08}, who showed that local galaxies exhibit an {\\it sublinear} SFR-HCN (J=3-2) relation, and thus follow the trend of decreasing SFR-HCN$^\\alpha$ index with increasing transition number characteristic of these models \\citep[e.g. Figure~\\ref{figure:sfr_lmol}; see also Figure 7 of ][]{nar08b}.  Similarly, this model satisfies the multi-line constraints of local CO observations.  The SFR-CO (J=1-0) relation in local galaxies appears to have an index ranging from $\\sim 1.3-1.5$. At higher-lying transitions (e.g. CO J=3-2), the index drops to $\\sim 0.9$, in accordance with theoretical predictions \\citep{ion09}.  We note that at lower bolometric luminosities, even the SFR-CO (J=1-0) relation may be subject to differential excitation and serve as a relatively poor tracer of the underlying Schmidt relation.      Finally, with an eye toward ALMA, we comment on the role of spatial resolution in observational determinations of the KS relation at high-\\z.  The exact mapping between the observed \\sigmaco \\ index, $\\alpha$, and the \\schmidt \\ index, $N$, depends on the level of thermalisation of the gas within the beam.  It is dependent (to first order) on the mean gas density.  Higher resolution observations which probe just the nucleus of the galaxy will probe higher mean densities, and allow even higher-lying CO emission lines to directly trace the underlying Schmidt SFR relation.  Indeed, some very tentative observational evidence for this trend in local galaxies has been shown by \\citet{nar08d}. The trends shown in Figures~\\ref{figure:sfr_lmol} and \\ref{figure:ks_molslope} are for the central 6 kpc of the galaxy, comparable to the typical resolution of current interferometric observations of \\zsim 2 galaxies \\citep[e.g. ][]{tac10}. Our models suggest that observations of the central $\\sim$2 kpc will probe sufficiently dense gas that the observed \\sigmacojthree \\ index, $\\alpha$, will trace the underlying \\schmidt \\ index, $N$."
    },
    "1005/1005.3811_arXiv.txt": {
        "abstract": "{We present wide-field {\\em Herschel}/PACS observations of A\\,1689, a massive galaxy cluster at $z{=}0.1832$, from our Open Time Key Programme.  We detect 39 spectroscopically confirmed 100$\\mu$m-selected cluster members down to $1.5{\\times}10^{10}L_{\\odot}$.  These galaxies are forming stars at rates in the range 1--10\\,M$_{\\odot}$/yr, and appear to comprise two distinct populations: two-thirds are unremarkable blue, late-type spirals found throughout the cluster; the remainder are dusty red sequence galaxies whose star formation is heavily obscured with A$({\\rm H}\\alpha){\\sim}2$\\,mag and are found only in the cluster outskirts. The specific-SFRs of these dusty red galaxies are lower than the blue late-types, suggesting that the former are in the process of being quenched, perhaps via pre-processing, the unobscured star formation being terminated first.  We also detect an excess of 100$\\mu$m-selected galaxies extending ${\\sim}6$Mpc in length along an axis that runs NE-SW through the cluster center at ${\\ga}9$5\\% confidence. Qualitatively this structure is consistent with previous reports of substructure in X-ray, lensing, and near-infrared maps of this cluster, further supporting the view that this cluster is a dynamically active, merging system.  } ",
        "introduction": "Over the last 30 years a large body of evidence has formed that star formation in cluster galaxies is much lower than in the surrounding field (e.g. Lewis et al. \\cite{lewis}).  A variety of mechanisms have been proposed to quench star formation in infalling spirals such as ram-pressure stripping or harassment (see Haines et al. \\cite{haines07} for a review), but the key distinguishing evidence remains elusive.  However most of these studies rely on optical spectroscopy using H$\\alpha$ or the [O{\\sc ii}] emission lines (e.g. Poggianti et al. \\cite{pog99}).  Recent mid-infrared based studies with {\\em ISO} and then {\\em Spitzer} (e.g. Metcalfe et al. \\cite{metcalfe}; Coia et al. \\cite{coia}; Geach et al. \\cite{geach06}; Saintonge et al. \\cite{saintonge}) have revealed a significant population of dusty cluster galaxies whose star formation is heavily obscured at optical wavelengths, leading to [O{\\sc ii}]-based SFRs underestimating the true level by ${\\sim}1$0--3$0{\\times}$ (Duc et al. \\cite{duc}), while the age-dependence of the dust obscuration can lead to dusty star-bursts being in fact classified spectroscopically as post-starburst galaxies (Poggianti \\& Wu \\cite{poggianti}).  Clusters are not isolated systems, but lie at the intersections of filaments, constantly accreting galaxies and galaxy groups along these filaments or directly from the field, and various studies suggest that the key sites of galaxy transformation are within these infalling structures rather than the cluster core (e.g.  Balogh et al. \\cite{balogh04}; Moran et al. \\cite{moran}; Fadda et al. \\cite{fadda08}).  These issues have motivated the Local Cluster Substructure (LoCuSS\\footnote{http://www.sr.bham.ac.uk/locuss}) \\emph{Herschel} Key Programme to study a large statistical sample of 30 massive galaxy clusters at $z{\\sim}0.2$ with panoramic FUV--FIR data from {\\em GALEX}, {\\em Spitzer} and {\\em Herschel} (Haines et al. \\cite{haines09a}; Smith et al. \\cite{smith10}).  These provide a complete census of star formation in cluster galaxies extending to the infall regions, which can then be related to the dynamical status of the cluster as determined from complementary lensing, X-ray and dynamical analyses (Haines et al. \\cite{haines09b}; Pereira et al. \\cite{pereira}).  \\begin{figure*} \\centering \\includegraphics*[height=101mm]{Herschelmap.eps}\\includegraphics[height=101mm]{A1689_merged.eps} \\caption{ ({\\em left:}) The {\\em Herschel}/PACS 100$\\mu$m map, with the spectroscopically confirmed cluster members indicated by blue open squares (scaling as 100$\\mu$m flux). Red contours indicate the projected $K$-band luminosity density of confirmed cluster members (weighted to account for spectroscopic incompleteness) estimated using the adaptive kernel method. ({\\em top-right}) Optical $g{-}i$ colour versus clustercentric radius. ({\\em bottom-right}) Redshift versus projected clustercentric radius. Blue, green and red symbols respectively indicate $K{\\le}K^{*}{+}2$ galaxies detected by: both {\\em Herschel} and {\\em Spitzer}; just {\\em Spitzer}; and neither telescope. Small black and grey points indicate low-mass ($K{>}K^{*}{+}2$) from ours and Czoske's spectroscopic surveys.  } \\label{cmradius} \\label{phasespace} \\label{map} \\end{figure*}  In this letter we present an analysis of the {\\em Herschel} (Pilbratt et al. \\cite{pilbratt}) FIR imaging of Abell 1689 at $z{=}0.1832$, one of the first clusters studied in the infrared with {\\em ISO} (Fadda et al. \\cite{fadda00}; Balogh et al. \\cite{balogh}, hereafter B02; Duc et al. \\cite{duc}, hereafter D02). Abell 1689 has many characteristics of a dynamically relaxed cluster, with a regular and highly symmetric X-ray morphology (Xue \\& Wu \\cite{xue}), the centroid of which coincides both with the brightest cluster galaxy and the centre of mass derived from strong lensing (Limousin et al. \\cite{limousin}).  However, the significant discrepancies obtained between mass estimates derived from strong lensing and X-ray analyses (Andersson \\& Madejski \\cite{andersson}), the very high concentration value observed for the cluster, and its large ($\\theta_{E}{\\sim}45^{\\prime\\prime}$) Einstein radius, all strongly favour a triaxial model in which the cluster is aligned along the line-of-sight (Oguri et al. \\cite{oguri}; Corless, King \\& Clowe \\cite{corless}).  Moreover, a substructure comprising ${\\sim}1$5\\% of the total cluster mass (Riemer-S{\\o}rensen et al. \\cite{riemer}) 1.5--2$^{\\prime}$ to the NE of the cluster centre has been identified from dynamical, X-ray and strong lensing analyses (Czoske \\cite{czoske}; Andersson \\& Madejski \\cite{andersson}; Limousin et al. \\cite{limousin}), while Kawaharada et al. (\\cite{kawaharada}, hereafter K10) find evidence for an accreting filament on 10--18\\,arcmin scales.  We present the data in \\S2, results in \\S3, and summary and discussion in \\S4. Throughout we assume a cosmology with $\\Omega_{M}{=}0.3$, $\\Omega_{\\Lambda}{=}0.7$ and H$_{0}{=}70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. ",
        "conclusions": "The results of the {\\em Herschel}/PACS photometry can be interpreted to suggest a filament of star-forming galaxies feeding the cluster A1689 at $z{=}0.1832$ from both sides. The direction of this filamentary structure coincides with the NE/SW axis of elongation of the cluster $K$-band light, and the observed asymmetry/substructure seen on 1--2\\,arcmin scales in previous X-ray and strong lensing analyses. K10 found anisotropies in gas temperature and entropy distributions in the cluster outskirts (10--18$^{\\prime}$), with a high $T_{X}{\\sim}5.4$\\,keV seen in the NE direction consistent with accretion flow along a filament, but low gas temperature (${\\sim}1.7$\\,keV) in all other directions. They also found evidence for a connecting large-scale structure (${\\sim}2$\\,deg) in the same NE direction in the form of SDSS galaxies with photometric redshifts consistent with the cluster.  The 100$\\mu$m selection has made the filaments stand out clearly against the background and dense cluster region, in the same way as seen by Fadda et al. (\\cite{fadda08}) for 24$\\mu$m sources in Abell 1763, due to the enhanced star-forming activity among the filament galaxies.  The SFRs of the {\\em Herschel} sources in A1689 are in the range 1--10\\,M$_{\\odot}{\\rm yr}^{-1}$, much lower than the 20--40\\,M$_{\\odot}{\\rm yr}^{-1}$ seen by Geach et al. (\\cite{geach06}, \\cite{geach}) for LIRGs in the rich cluster Cl\\,0024+16 at $z{=}0.4$. If this star formation is to build the bulges needed to turn the spirals infalling into clusters at $z{\\sim}0.4$ into the S0s found in local clusters it would require ${\\ga}$1--3\\,Gyr at the SFRs observed here and in other $z{\\sim}0.2$ clusters (Smith et al. \\cite{smith10}), rather than the ${\\sim}10^{8}$yr implied by Geach et al. (\\cite{geach})'s results.   We find two populations of {\\em Herschel} sources in A1689, analogous to the blue cloud and dusty-red galaxies identified by Wolf et al. (\\cite{wolf05},\\cite{wolf09}) in A901/2, and we find similar EW(O{\\sc ii}), EW(H$\\delta$) and specific-SFRs to them.  In particular, for the dusty red galaxies we obtain $\\langle$\\,EW(O{\\sc ii})$\\rangle{\\sim}5${\\AA}, due to having A(O{\\sc ii}$){\\sim}4$\\,mag placing them on the k+a/e(a) boundary, yet their specific-SFRs are systematically ${\\sim}3{\\times}$ lower than those of the blue cloud population (Fig.~\\ref{ssfr}), suggesting that they are currently being quenched. Whereas the ``blue cloud'' population are found throughout the cluster, the ``dusty red'' galaxies are found only outside the cluster core ($r{\\ga}1$\\,Mpc), where they make up half of the star-forming galaxies detected by {\\em Herschel} (Fig.~\\ref{cmradius}; top-right). This is consistent with the findings of STAGES that the dusty red galaxies prefer the intermediate-density environment of cluster infall regions (Gallazzi et al. \\cite{gallazzi}; Wolf et al. \\cite{wolf09}), and morphologically they found ${\\sim}70$\\% to be Sa/Sb spirals with a bright nucleus or inner bar/disk, as well as a few cases of interactions or merger remnants, suggestive of galaxy harassment, interactions or mergers for their origin.  Four of the brightest eight {\\em Herschel} detections appear associated with groups: two are within a compact system of 3 galaxies ($\\sigma_{\\nu}{\\sim}120$\\,km\\,s$^{-1}$; labelled A in Fig.~\\ref{map}); one within a $\\sigma_{\\nu}{\\sim}320$\\,km\\,s$^{-1}$ group (16 members; group B); and the last within the substructure identified by Czoske (\\cite{czoske}; group C). These, plus similar results obtained by Pereira et al. (\\cite{pereira}) for Abell 1835, suggest that merger-induced starbursts and transformations in low-mass groups (pre-processing) could represent a significant contribution to the high-SFR cluster galaxy population, and an important evolutionary pathway for the formation of S0s.  The availability of FUV--FIR imaging plus extensive spectroscopy for a sample of 30 clusters spanning the full range of morphologies and dynamical states within LoCuSS will allow us in future to estimate the relative fractions of infalling galaxies accreted by clusters from filaments, groups and the field, and measure the importance of pre-processing inside groups.  \\begin{figure} \\centering \\includegraphics[width=68mm]{SSFR_hist.eps} \\caption{Histograms of specific-SFRs for 100$\\mu$m-detected cluster members within the red sequence (orange stripes) and blue cloud (blue).} \\label{ssfr} \\end{figure}"
    },
    "1005/1005.0509_arXiv.txt": {
        "abstract": "{Star formation rate (SFR), metallicity and stellar mass are within the important parameters of star--forming (SF) galaxies that characterize their formation and evolution. They are known to be related to each other at low and high redshift  in the mass--metallicity, mass--SFR, and metallicity--SFR relations.} % {In this work we demonstrate the existence of a plane in the 3D space defined by the axes SFR [log(SFR)(M$_{\\odot}$yr$^{-1}$)], gas metallicity [12+log(O/H)], and stellar mass [log(M$_{\\rm star}$/M$_{\\odot}$)] of SF galaxies.} {We used SF galaxies from the $``$main galaxy sample$\"$ of the Sloan Digital Sky Survey--Data Release 7 (SDSS --DR7) in the redshift range $0.04 < z < 0.1$ and $r-$magnitudes between 14.5 and 17.77. Metallicities, SFRs, and stellar masses were taken from the Max-Planck-Institute for Astrophysics--John Hopkins University (MPA-JHU) emission line analysis database.} {From a final sample of 32575 galaxies, we find for the first time a fundamental plane for field galaxies relating the SFR, gas metallicity, and stellar mass for SF galaxies in the local universe. One of the applications of this plane would be estimating stellar masses from SFR and metallicity. High redshift data from the literature at redshift $\\sim$0.85, 2.2, and 3.5, do not show evidence for evolution in this fundamental plane.} {} ",
        "introduction": "Relations between important properties of astrophysical objects often lead to the discovery of the so--called fundamental planes when three parameters are involved. The fundamental plane (FP) for elliptical galaxies \\citep{Djor87,Dressler87}, relates their luminosity, velocity dispersion (dynamics), and scale length (morphology). This FP represents an important tool to investigate the properties of early type and dwarf galaxies, to perform cosmological tests, and compute cosmological parameters. It is also an important diagnostic tool for galaxy evolution and mass--to--light ($M/L$) variations with redshift.  Fundamental planes have also been defined for  globular clusters  \\citep{Meylan97} and galaxy clusters \\citep{Schaeffer93, Adami98}. The parameter space of globular clusters, elliptical galaxies, and galaxy clusters is properly described by a geometrical plane $L \\propto R^{\\alpha}\\sigma^{\\beta} $, where $L$ is the optical luminosity of the system, $R$ is a measure of the size scale, $\\sigma$ is the velocity dispersion of the system, and $\\alpha$ and $\\beta$ are free parameters. The FP for globular clusters, elliptical galaxies, and galaxy clusters have very similar slopes, which means that, accounting for differences in zero points, a single FP with a range of about nine orders of magnitude in luminosity can be defined \\citep{Schaeffer93,Ibarra09}. %  The FP that we introduce here relates three fundamental parameters: the SFR [log(SFR)(M$_{\\odot}$yr$^{-1}$)], gas metallicity [12+log(O/H)], and stellar mass [log(M$_{\\rm star}$/M$_{\\odot}$)] of field SF galaxies. All these variables have been related in the past by the mass--metallicity ($M-Z$) relation  \\citep{Lequeux79}, the mass--SFR relation \\citep{Brinchmann04}, and the metallicity--SFR relation \\citep{Lara10,Lopez-Sanchez10}. Also, some authors have studied the inter-dependence of those variables \\citep[e.g.][]{Hoopes07, Ellison08, Mannucci10}. In the present work, we propose the generalization of those relations defining a plane formed by a linear combination of two of those variables with respect to the third one.  The $M-Z$ relates the mass and metallicity of galaxies, with massive galaxies showing higher metallicities than less massive ones, and it has been well established for the local universe ($z \\sim$ 0.1) by the work of \\citet{Tremonti04} using SDSS data. The $M-Z$ relation has also been studied at low redshifts $z \\sim 0.35$ \\citep{Lara09a,Lara09b}, at intermediate redshifts $z \\sim 0.7$ \\citep[e.g.,][]{Rodrigues08}, and at high redshift $z \\sim 2.2$ and $z \\sim 3.5$ \\citep[][respectively]{Erb06,Maiolino08}.  The stellar mass of SF galaxies is also related to the SFR, in the sense that more massive galaxies show higher SFRs \\citep{Brinchmann04,Salim05}. However, \\citet{Brinchmann04} emphasized that at log(M$_{\\rm star}$/M$_{\\odot}$) $\\gtrsim$ 10, the distribution of SFRs broadens significantly and the correlation between stellar mass and SFR breaks down. At higher redshifts, \\citet{Noeske07} showed the existence of a $``$main sequence$\"$ (MS) for this relation over the redshift range $0.2 < z < 1.1$, with the slope of the MS  moving as a whole as $z$ increases.  The metallicity and SFR of SF galaxies are weakly correlated, as will be observed in Fig. 1. However, and despite of the high scatter, SFR increases with metallicity \\citep{Lara10,Lopez-Sanchez10}.  This paper is structured as follows: Sect. 2  describes the data selection as well as the SFRs, metallicities and stellar masses estimations given by the Max-Planck-Institute for Astrophysics--John Hopkins University (MPA-JHU) group\\footnote{http://www.mpa-garching.mpg.de/SDSS} and adopted in this work. In Sect. 3 we define the FP for field galaxies, and conclusions are given in Sect. 4. ",
        "conclusions": "We have demonstrated the existence of a FP for field SF galaxies in the 3D space formed by the orthogonal coordinate axes log(M$_{\\rm star}$/ M$_{\\odot}$), log(SFR)(M$_{\\odot}\\rm yr^{-1}$), and 12+log(O/H); three of the fundamental parameters of galaxies. All those variables have been related previously in pairs as with the $M-Z$, metallicity--SFR, and mass-SFR relations, but this is the first time that the correlation for all of them has been quantified.  The FP presented here allows estimating the stellar mass [log(M$_{\\rm star}$/ M$_{\\odot}$)] of field galaxies as a linear combination of 12+log(O/H) and log(SFR)(M$_{\\odot}\\rm yr^{-1}$). The scatter in the mass estimates using the FP (1$\\sigma$ error of 0.16) is lower that that obtained through the $M-Z$ and mass-SFR relations.  The FP introduced here would be useful for deriving masses in spectroscopic surveys where the SFR and metallicity are estimated for emission line galaxies, for example, using the H$\\alpha$ luminosity to estimate the SFR \\citep[e.g.][]{Kennicutt98}, and any of the metallicity methods in the literature, such as the $R_{23}$ \\citep{Pagel79} or N2 \\citep{Denicolo02}, see \\citet{Kewley08} for a review. However, since this study has been carried out using emission line galaxies, this FP will be useful only when both, SFR and metallicity of galaxies can be estimated.  Within the errors, there is no evidence of an evolution of the local FP when applied to high redshift samples. Which means that it could be useful even at high redshifts, where measuring the continuum and absorption lines for fitting models would be more difficult and time consuming.  Then, we propose the use of this FP as an alternative tool to the existing methods to determine the stellar mass of galaxies at low and high redshifts."
    },
    "1005/1005.5699_arXiv.txt": {
        "abstract": "{The standard cooling models of neutron stars predict temperatures of $T<10^{4}$~K for ages $t>10^{7}$~yr. However, the likely thermal emission detected from the  millisecond pulsar J0437-4715, of spin-down age $t_s \\sim 7\\times10^9$~yr, implies a temperature  $T\\sim 10^5$~K. Thus, a heating mechanism needs to be added to the cooling models in order to obtain agreement between theory and observation.} {Several internal heating mechanisms could be operating in neutron stars, such as magnetic field decay, dark matter accretion, crust cracking, superfluid vortex creep, and non-equilibrium reactions  (``rotochemical heating''). We study these mechanisms to establish which could be the dominant source of thermal emission from old pulsars.} {We show by simple estimates that magnetic field decay, dark matter accretion, and crust cracking are unlikely\tto have a significant heating effect on old neutron stars. The thermal evolution for the other mechanisms is computed with the code of Fern\\'andez and Reisenegger. Given the dependence  of the heating mechanisms on the spin-down parameters, we study the thermal evolution for two types of pulsars: young, slowly rotating  ``classical'' pulsars  and old, fast rotating millisecond pulsars.} {We find that magnetic field decay, dark matter accretion, and crust cracking do not produce any detectable heating of old pulsars. Rotochemical heating and vortex creep can be important both for classical pulsars and millisecond pulsars. More restrictive upper limits on the surface temperatures of classical pulsars could rule out  vortex creep as the main source of thermal emission. Rotochemical heating in classical pulsars is driven by the chemical imbalance built up during their early spin-down, and is therefore strongly sensitive to their initial rotation period.} {} ",
        "introduction": "\\label{sec:intro} Neutron stars (NSs) are compact objects composed of a  liquid core enveloped by a solid crust. The core is expected to contain superfluid neutrons and superconducting protons, while the crust contains heavy atomic nuclei arranged in a crystal lattice,  coexisting with superfluid neutrons in its inner part. The high density of the NS core, up to $\\sim(3-9)\\rho_0$, where $\\rho_0$ is  the saturation nuclear matter density, cannot be reproduced in terrestrial laboratories. This turns the NSs into natural laboratories. The study of their thermal evolution, confronting theory and observation, provides a useful test for the understanding of the properties of matter at supernuclear density.  For all standard cooling models \\citep{yak04}, neutron stars cool down to surface temperatures $T_s < 10^4$~K within less than $10^7$~yr. Nevertheless, the observation of ultraviolet thermal emission from millisecond pulsar J0437-4715 \\citep{kar04}, whose spin-down age, corrected to the latest distance of $157$~pc \\citep{del08}, is $\\tau_{sd}\\sim 7\\times 10^9$~yr \\citep{van01}, shows a surface temperature of about $\\sim 10^5$~K for this pulsar. Hence, a heating mechanism needs be to added  to the standard cooling models  to obtain agreement between theory and observation.   There are several heating mechanisms that can be present during the late stages of the thermal evolution. These include the frictional motion of superfluid neutron vortices \\citep{alp84,shi89,lar99}, rotochemical heating \\citep{reis95,reis97,fer05,petro09}, magnetic field decay \\citep{gol92,thom96,pon07}, and crust cracking \\citep{baym71,chen92}. Other mechanisms, based on more speculative hypotheses, such as a time variation of the gravitational constant \\citep{jof06}, the decay of exotic particles \\citep{han02}, or the accretion of dark matter particles \\citep{dela10,kouv10}, could in principle also heat old neutron stars. In \\citet{sch99} and \\citet{lar99}, several of the internal heating mechanisms cited are studied and confronted with observational data of neutron stars (surface temperature upper limits). However, the oversimplified description of some heating mechanisms \\citep[rotochemical heating in][]{sch99}  and the until then non-detection of thermal emission from neutron stars older than $\\sim10^6$ yr, made it impossible to obtain reliable conclusions.  The goal of this work is  to provide a comparative analysis of the thermal evolution including different heating  mechanisms. In order to do this, we discard some of them as not strong enough (magnetic field decay, dark matter accretion, and crust cracking), and we present a more detailed study of the most promising ones: vortex creep and rotochemical heating. We confront these mechanisms with  the thermal emission detected in the millisecond pulsar J0437-4715 and the best available upper limits  on the temperature of  six other old pulsars. Owing to the dependence of heating mechanisms on spin-down parameters, which leads to different temperatures for different pulsars, we separately study the thermal evolution for two types of pulsars: young, strongly magnetized, and slowly rotating ``classical'' pulsars, and old, weakly magnetic, and fast rotating millisecond pulsars (MSPs).  The paper is organized as follows. In Sect.~\\ref{sec:mech} we describe the different heating processes and assess their importance. In Sect.~\\ref{sec:eff} we show the effects of vortex creep and rotochemical heating in classical and millisecond pulsars. In Sect.~\\ref{sec:obs} the predictions are confronted with observed data. A summary of our main conclusions is given in Sect.~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We studied five heating mechanisms that can be operating in old neutron stars: magnetic field decay, dark matter accretion, crust cracking, vortex creep, and rotochemical heating, and compared them with pulsar observations.  We found that magnetic field decay, dark matter accretion, and crust cracking cannot produce detectable heating in old pulsars. Owing to the high yield strain angle \\citep{hor09}, the crust cracking mechanism does not operate in classical pulsars, and probably only operates in MSPs. The vortex creep and rotochemical heating can be important both for classical pulsars and MSPs.  In the evolutionary curves with vortex creep, and with the excess angular  momentum, $J$, set to the lowest temperature compatible with the thermal emission detected from MSP J0437-4715, the predicted  temperature turns out to be very near the observational upper limits for the classical pulsars B1929+10 and B1133+16. Likewise, rotochemical heating with modified Urca reactions and no Cooper pairing is only 1.7 $\\sigma$ below the temperature measured in the MSP J0437-4715. However, the temperature prediction of rotochemical heating can be raised if a superfluid core is considered in the model.  The prediction of pinning energies in the inner crust of NSs via the semi-classical model for the vortex-nuclei interaction \\citep{don04} overestimates the temperatures  of B0950+08 and B1929+10. The recent estimations of pinning energies via a quantum approach \\citep{avo08} are consistent with the observations of classical pulsars and MSPs. Finally, more stringent constraints on the temperature of some classical pulsars such as B0950+08 could rule out the vortex creep mechanism as the main source of the thermal emission detected in the MSP J0437-4715."
    },
    "1005/1005.0844_arXiv.txt": {
        "abstract": "The growth of the supermassive black holes (BHs) that reside at the centres of most galaxies is intertwined with the physical processes that drive the formation of the galaxies themselves.  The evolution of the relations between the mass of the BH, $\\mbh$, and the properties of its host therefore represent crucial aspects of the galaxy formation process.  We use a cosmological simulation, as well as an analytical model, to investigate how and why the scaling relations for BHs evolve with cosmic time.  We find that a simulation that reproduces the observed redshift zero relations between $\\mbh$ and the properties of its host galaxy, as well as the thermodynamic profiles of the intragroup medium, also reproduces the observed evolution in the ratio $\\mbh/\\ms$ for massive galaxies, although the evolution of the $\\mbh/\\sigma$ relation is in apparent conflict with observations. The simulation predicts that the relations between $\\mbh$ and the binding energies of both the galaxy and its dark matter halo do not evolve, while the ratio $\\mbh/\\mhalo$ increases with redshift. The simple, analytic model of \\citet{boot10}, in which the mass of the BH is controlled by the gravitational binding energy of its host halo, quantitatively reproduces the latter two results. Finally, we can explain the evolution in the relations between $\\mbh$ and the mass and binding energy of the stellar component of its host galaxy for massive galaxies ($\\ms\\sim 10^{11}\\,\\msun$) at low redshift ($z<1$) if these galaxies grow primarily through dry mergers. ",
        "introduction": "\\label{sec:intro}  Over the past decade it has become clear that the supermassive black holes (BHs) found at the centres of virtually all galaxies with spheroidal components, have masses that are coupled to the properties of their host galaxies \\citep{magg98,ferr00,trem02,hari04,hopk07b}. Additionally, there exists evidence that BH masses are coupled to the properties of the dark matter haloes in which they reside \\citep{ferr02,boot10}. Further correlations between quasar activity \\citep[e.g.][]{boyl98} and the evolution of the cosmic star formation rate \\citep[e.g.][]{mada96} provide evidence that there exists a link between galactic star formation and accretion onto a central AGN.  It has long been recognised that the growth of BHs is likely self-regulated \\citep{silk98} and that these tight correlations indicate that the growth of BHs is tightly intertwined with the physical processes that drive galaxy formation.  However, despite a wide variety of theoretical and observational studies, the origin of these relations is still debated. The study of the evolution of the BH scaling relations therefore represents a crucial aspect of the galaxy formation process that may provide us with additional clues regarding the physical processes that give rise to the BH scaling relations.  Addressing these questions observationally is challenging. Due to their extremely high luminosities, bright quasars provide a promising route to measuring BH masses at high redshift through the widths of low-ionization lines that are associated with the broad-line region close to the BH and using the assumption of virial equilibrium \\citep[e.g.][]{vest02}.  It has, however, been claimed that this procedure systematically underestimates BH masses \\citep{jarv02}. Measuring galaxy masses for these objects is very difficult as the BH outshines the galaxy by a large factor \\citep[see e.g.\\ the discussion  in ][]{merl09}. Since AGN surveys are biased towards more massive black holes, selection effects also need to be taken into account \\citep[e.g.][]{shen09,benn10}, which can make it difficult to distinguish between evolution in the normalization and in the scatter in the scaling relations \\citep{laue07}. In spite of these difficulties, measurements of the BH scaling relations have been made as far out as redshift three.  \\citet{mclu06} found that the BHs associated with radio loud AGN residing in galaxies of a given stellar mass are a factor of four more massive at redshift two than in the local Universe.  \\citet{deca09} studied the \\civ\\, line associated with the quasar broad line region in R-band selected hosts at both redshifts zero and three and found that BHs are typically a factor of seven more massive at high redshift for a given galaxy mass.  These results are consistent with other observational studies \\citep{walt04,peng06,peng06a,merl09,gree10,benn10}. Taken together, these papers suggest an emerging consensus that at higher redshift BHs in hosts of a given mass are systematically more massive than in the local Universe, although see \\citet{jahn09} for one study that finds no significant evolution.  The evolution of the relation between BH mass, $\\mbh$, and stellar velocity dispersion, $\\sigmas$, has been studied utilising the width of the \\oiii\\, line as a proxy for stellar velocity dispersion \\citep{nels96}.  These studies suggest that the $\\mbh-\\sigmas$ relation either does not evolve \\citep{shie03,gask09}, or does so weakly, with BHs $\\sim 0.1-0.3$~dex more massive at $z=1$ \\citep{salv06,gu09,woo08,treu07}.  The evolution of the BH scaling relations has also been studied using numerical simulations \\citep[e.g.][]{robe06,joha09} and semi-analytic models \\citep[e.g.][]{malb07,lama10,kisa10}. \\citet{robe06} employed simulations of idealised galaxy mergers, initialised to have properties typical of merger progenitors at various redshifts, to construct the relation between galaxy stellar mass, $\\ms$, and $\\sigmas$ as a function of redshift and found that, at a given value of $\\sigmas$, the corresponding $\\mbh$ decreases mildly with increasing redshift.  At $z=1$ the simulations of \\citet{dima08} have BHs that lie slightly above the $z=0$ normalization of the $\\mbh-\\sigma$ relation.  However, these simulations were stopped at $z=1$ and so cannot inform us about the evolution of the $\\mbh-\\sigmas$ toward lower redshift.  However, for $z>1$ they predict a weak evolution in the $\\mbh-\\sigmas$ relation such that at higher redshift galaxies of a given velocity dispersion contain slightly less massive BHs.  \\citet{joha09} employed similar numerical techniques to argue that it is unlikely that BHs are able to form significantly before their host bulges. Semi-analytic models that reproduce many redshift zero properties of galaxies also predict that, at a fixed $\\sigmas$, BH masses decrease with increasing redshift \\citep{malb07}. These theoretical models thus predict evolutionary trends that go in the opposite direction to those inferred from observations.  Finally, the models of \\citet{hopk09} predict that, at a fixed stellar velocity dispersion, BH masses at higher redshift are either the same (for $\\mbh\\sim10^8\\,\\msun$) or slightly more massive (for $\\mbh>10^8\\,\\msun$) at fixed $\\sigmas$ than their redshift zero counterparts, in agreement with observation.  On the other hand, the relation between BH mass and galaxy bulge mass shows a positive evolution in both semi-analytic models \\citep{malb07, hopk09} and numerical simulations \\citep{dima08}, the magnitude of which is comparable to that observed.  The larger spread in the predictions for the evolution of the $\\mbh-\\sigmas$ relation may reflect that it is more difficult to predict velocity dispersions, which depend on both mass and size, than it is to predict masses.  In \\citet[][hereafter BS09]{boot09} we presented self-consistent, hydrodynamical simulations of the co-evolution of the BH and galaxy populations that reproduce the redshift zero BH scaling relations. These same simulations also match group temperature, entropy and metallicity profiles, as well as the stellar masses and age distributions of brightest group galaxies \\citep{mcca10}.  In \\citet[]{boot10} (hereafter BS10) we used the same simulations, as well as an analytic model, to demonstrate that $\\mbh$ is determined by the mass of the dark matter (DM) halo with a secondary dependence on the halo concentration, of the form that would be expected if the halo binding energy were the fundamental property that controls the mass of the BH.  In the present work we use the same models to investigate why and how the BH scaling relations evolve for massive galaxies.  This paper is organised as follows. In Sec.~\\ref{sec:method} we summarise the numerical methods employed in this study and the simulation analysed.  In Sec.~\\ref{sec:ev} we present predictions for the evolution of the BH scaling relations and compare them to observations. We find that the evolution in the $\\mbh-\\ms$ relation predicted by the simulations is in excellent agreement with the observations, while the measured weak evolution in the $\\mbh-\\sigma$ relation is in apparent disagreement, and predict that while BH mass increases with redshift for fixed halo mass, the relations between $\\mbh$ and the binding energies of both the host galaxies and DM haloes do not evolve. We demonstrate in \\ref{sec:explanation} that the a analytic description in which $\\mbh$ is coupled to the DM halo binding energy can reproduce the evolution of the relation between BH and halo mass. Furthermore, we show that the evolution in the relations between the BH and the stellar mass and binding energy can be understood in terms of the more fundamental relation with the binding energy of the dark halo and the growth of massive galaxies through dry mergers.  Finally, we summarise our main conclusions in Sec.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have used a self-consistent cosmological simulation that reproduces the observed redshift zero relations between $\\mbh$ and both galaxy and halo properties (as well as the thermodynamic profiles of the intragroup medium) to investigate how, and why, these relations evolve through time.  The relation between BH mass and host galaxy mass predicted by the simulation is consistent with available observations at $z<2$, which are currently confined to $\\ms\\sim 10^{11}\\,\\msun$ galaxies. For such galaxies we predict that the ratio $\\mbh/\\ms \\propto (1+z)^{\\alphas}$, with $\\alphas \\approx 0.5$, and $\\mbh/\\sigmas^4 \\propto (1+z)^{\\alphas}$ with $\\alphas\\approx-0.3$, in apparent conflict with recent observations.  The ratio between the BH mass and the binding energy of the dark halo is independent of redshift, in agreement with BS10 who argued that the BH mass is controlled by the halo binding energy.  The simple analytic model of BS10, in which the BH mass is assumed to scale in proportion to the binding energy of the dark halo, not only reproduces the simulated redshift zero $\\mbh-\\mhalo$ relation, but also its evolution. For a fixed halo mass BHs are more massive at higher redshift because the haloes are more compact and thus more tightly bound. Assuming an NFW halo density profile and the evolution of the halo concentration-mass relation predicted by simulations, the model can quantitatively account for the predicted evolution.  The simulation predicts that the ratio between the BH mass and the binding energy of the stellar component of the galaxy is also independent of redshift (at least for $\\ms\\sim10^{11}\\,\\msun$ and $z<1$), even though BS10 demonstrated explicitly that the correlations between BH mass and stellar properties are not fundamental. This result is, however, consistent with a picture in which massive galaxies grow primarily through dry mergers at low redshift, which we showed to be the case in the simulation. Combined with the observed evolution in the $\\ms-\\sigmas$ relation, this idea can quantitatively account for the evolution in the $\\mbh-\\ms$ relation.  One interesting implication of this scenario is that the evolution of the relations between BHs and the properties of their host galaxies may differ for galaxies that do not grow predominantly through dry mergers, as would be expected for lower masses and at higher redshifts. We will investigate this further in the a future work, employing higher resolution simulations."
    },
    "1005/1005.2888_arXiv.txt": {
        "abstract": "Evidence is presented indicating that in the hard state of Cygnus X-1, the coronal magnetic field might be below equipartition with radiation (suggesting that the corona is not powered by magnetic field dissipation) and that the ion temperature in the corona is significantly lower than what predicted by ADAF like models.  It is also shown that the current estimates of the jet power set interesting contraints on the jet velocity (which is at least mildly relativistic), the accretion efficiency (which is large in both spectral states), and the nature of the X-ray emitting region (which is unlikely to be the jet). ",
        "introduction": "\\label{sec:intro}   Black hole binaries are observed in two main X-ray spectral states, namely the Hard State (HS) and the Soft State (SS), see [1]. The hard X-ray emission in both spectral states is well represented by Comptonisation by an hybrid thermal/non-thermal electron distribution. In the HS the temperature and optical depth of the thermal electrons are higher, and the slope of the non-thermal tail seem steeper than in the SS.  Consequently, the hard X-ray emission is dominated by thermal Comptonisation in the HS and by non-thermal Comptonisation in the SS.  The HS is known to be associated with the presence of a compact radio jet which   observed in the SS.  Here we focus  on the prototypical black hole source Cygnus X-1. In section~\\ref{sec:model}, we present a relatively simple coupled kinetic-radiation model that  allows us to understand the origin of the very different spectral shapes observed in the two spectral states as well as  the spectral evolution during state transitions (see e.g. [2]). A thorough investigation of the model and its discussion in the context of the observations  can be found  in [3]; the present paper summarises our main results. Then in section~\\ref{sec:jet}, we summarise the arguments developed in [4] showing that the present estimates of the jet power of Cygnus X-1 imply that the jet has a relativistic velocity, that the accretion proceeds efficiently in the HS and that the X-ray emission is unlikely to be produced in the jet. ",
        "conclusions": "In both spectral states of black hole binaries the coronal emission  can be powered by a similar non-thermal acceleration mechanism. In the HS the synchrotron and $e$-$e$ Coulomb boilers redistribute the energy of the non-thermal particles to form and keep a quasi-thermal electron distribution at a relatively high temperature, so that most of the luminosity is released through quasi-thermal Comptonisation. In the SS,  the soft photon flux from the accretion disc becomes very strong and cools down the electrons, reducing the thermal Compton emissivity. This change in the soft photon flux could be caused either because the inner radius of the truncated disc moves inward into the central hot accretion flow, or, in the framework of accretion disc corona models, because the disc temperature increases dramatically. Then most of the emission is produced by disc photons up-scattered by the non-thermal cooling electrons. Our comparison of simulations with the high energy spectra of Cygnus X-1 in the HS allowed us to set upper limits on the magnetic field and the proton temperature.  Our results indicate that, { as long as the non-thermal MeV tail is produced in the same region as the bulk of the Comptonised emission}, the  magnetic field  in the HS  is below equipatition with radiation (unlike what is assumed in most accretion disc corona models). The proton temperature is found to be significantly lower than predicted by standard 2-temperature accretion flow models ($kT_{\\rm i}<$2 MeV). We also note that such accretion flows are usually radiatively inneficient while the jet energetics suggests efficient accretion in the HS. The present estimates of the jet power also suggest that the jet is Thomson thin with a velocity which is {at least} mildly relativistic, which does not make it a favoured location for the production of the observed X-ray emission."
    },
    "1005/1005.3076_arXiv.txt": {
        "abstract": "{ In \\textit{Herschel} images of the Galactic plane and many star forming regions, a major factor limiting our ability to extract faint compact sources is cirrus confusion noise, operationally defined as the ``statistical error to be expected in photometric measurements due to confusion in a background of fluctuating surface brightness.''  The histogram of the flux densities of extracted sources shows a distinctive faint-end cutoff below which the catalog suffers from incompleteness and the flux densities become unreliable.  This empirical cutoff should be closely related to the estimated cirrus noise and we show that this is the case.  We compute the cirrus noise directly, both on \\textit{Herschel} images from which the bright sources have been removed and on simulated images of cirrus with statistically similar fluctuations.  We connect these direct estimates with those from power spectrum analysis, which has been used extensively to predict the cirrus noise and provides insight into how it depends on various statistical properties and photometric operational parameters. We report multi-wavelength power spectra of diffuse Galactic dust emission from Hi-GAL observations at 70 to 500\\,$\\mu m$ within Galactic plane fields at $l= 30 \\degr$ and $l= 59 \\degr$.  We find that the exponent of the power spectrum is about $-3$.  At $250\\ \\mu{\\rm m}$, the amplitude of the power spectrum increases roughly as the square of the median brightness of the map and so the expected cirrus noise scales linearly with the median brightness.  For a given region, the wavelength dependence of the amplitude can be described by the square of the spectral energy distribution (SED) of the dust emission. Generally, the confusion noise will be a worse problem at longer wavelengths, because of the combination of lower angular resolution and the rising power spectrum of cirrus toward lower spatial frequencies, but the photometric signal to noise will also depend on the relative SED of the source compared to the cirrus.} ",
        "introduction": "Cirrus noise, which is operationally defined as the ``statistical error to be expected in photometric measurements due to confusion in a background of fluctuating surface brightness'' \\citep{gautier}, is a major issue limiting the cataloging of compact sources that underpins the study of the early stages of star formation in the interstellar medium.  Examination of wide range of mass of the stellar precursors, requires measurement of sources with a wide range of luminosity, or at each wavelength, flux density. Stars form where there is abundant material, and so the cirrus brightness in the field tends to be high. Furthermore, many studies, both targeted and unbiased, are in the Galactic plane, which is also bright.  Cirrus noise varies with cirrus brightness (Sect.~\\ref{sect:cnps}), introducing further complexity to the problem. Cirrus fluctuations characteristically decrease with decreasing spatial scale (Sect.~\\ref{sect:cnps}), but even with the improved angular resolution of \\textit{Herschel}, cirrus noise remains a dominant factor. The \\textit{Herschel} observation planning tool HSpot (www.ipac.caltech.edu/Herschel/hspot.shtml) has a built-in confusion noise estimator to provide ``on-line guidance on where to expect fundamental detection limits for point sources that cannot be improved by increasing the integration time.'' With such guidance the serious impact of cirrus noise was anticipated, and as shown below it can now be quantified directly using submillimeter data.  To appreciate the problem at its fundamental level, consider actual catalogs of extracted sources \\citep{MolinariThisVolume,mol2010detect} in two degree-sized Hi-GAL Galactic plane fields \\citep{molinari2010} (see Sect.~\\ref{sect:obs}). \\begin{figure}% \\centering \\includegraphics[width=0.4\\textwidth, angle=0]{14684fg1.ps} \\caption{ Solid histogram: 136 sources cataloged at 250\\,$\\mu m$ in a representative large degree-sized sub-region of Hi-GAL $l= 30\\degr$ field (f30: see Sect.~\\ref{sect:obs}).  Dashed histogram: 70 sources cataloged in a sub-region of Hi-GAL $l= 59\\degr$ field (f59) that has a factor 5.8 lower median brightness.} \\label{fig:hist2} \\end{figure} The solid histogram in Figure~\\ref{fig:hist2} shows the flux densities of sources cataloged at 250\\,$\\mu m$ in the brighter field.  Note the falloff in source counts at flux densities less than 10~Jy. This falloff is not an intrinsic property of the underlying population of sources. We show below that this falloff is the expected consequence of cirrus noise, which can be quantified independently of the making of the catalogs.  In the fainter field (dashed histogram), the cutoff is lower because of decreased cirrus noise. Clearly, it will be important to account for the varying dramatic effects of cirrus noise in order to recover the statistics of the intrinsic faint-source population.%  Cirrus noise is operationally defined and so for a particular measuring strategy, such as fitting compact sources with Gaussians as used in Hi-GAL, this can actually be estimated directly from the source-removed maps (Sect.~\\ref{sect:cnps}). There is also an extensive literature on ``estimating confusion noise due to extended structures given some statistical properties of the sky'' \\citep{gautier}; see \\citep{kiss-galcir,jeong2005,ma-dust,roy2010}.  The cirrus brightness statistics are well described by a power spectrum, appropriate for Gaussian random fields.  \\citet{gautier} found some non-Gaussianity and \\citet{ma-dust} found non-vanishing skewness and excess kurtosis in the underlying brightness fluctuation fields. Nevertheless, these are not large effects and for estimating the variance the power spectrum is demonstrably still a powerful tool, particularly because of the insight it provides into how the cirrus noise depends on various statistical and operational parameters. Therefore, we connect the new direct estimates with what can be obtained using power spectra (Sect.~\\ref{sect:cnps}).  In Sect.~\\ref{sect:prop} we show that reliable power spectra can be obtained even from first-generation images processed for Hi-GAL. Finally, we return to how estimates of cirrus noise should ultimately explain the faint-end cutoff in forthcoming source catalogs (Sect.~\\ref{sect:noise}). ",
        "conclusions": ""
    },
    "1005/1005.1073_arXiv.txt": {
        "abstract": "{ The constituents of the cosmic IR background (CIB) are studied at its peak wavelengths (100 and 160 $\\mu$m) by exploiting Herschel/PACS observations of the GOODS-N, Lockman Hole, and COSMOS fields in the {\\em PACS Evolutionary Probe} (PEP) guaranteed-time survey. The GOODS-N data reach 3$\\sigma$ depths of $\\sim$3.0 mJy at 100 $\\mu$m and $\\sim$5.7 mJy at 160 $\\mu$m. At these levels, source densities are 40 and 18 beams/source, respectively, thus hitting the confusion limit at 160$\\mu$m. Differential number counts extend from a few mJy up to 100-200 mJy, and are approximated as a double power law, with the break lying between 5 and 10 mJy. The available ancillary information allows us to split number counts into redshift bins. At $z\\le0.5$ we isolate a class of luminous sources ($L_{IR}\\sim10^{11}$ L$_\\odot$), whose SEDs resemble late-spiral galaxies, peaking at $\\sim$130 $\\mu$m restframe and significantly colder than what is expected on the basis of pre-Herschel models. By integrating number counts over the whole covered flux range, we obtain a surface brightness of $6.36\\pm1.67$ and $6.58\\pm1.62$ $[$nW m$^{-2}$ sr$^{-1}]$ at 100 and 160 $\\mu$m, resolving $\\sim$45\\% and $\\sim$52\\% of the CIB, respectively. When stacking 24 $\\mu$m sources, the inferred CIB lies within 1.1$\\sigma$ and 0.5$\\sigma$ from direct measurements in the two bands, and fractions increase to 50\\% and 75\\%. Most of this resolved CIB fraction was radiated at $z\\le1.0$, with 160 $\\mu$m sources found at higher redshift than 100 $\\mu$m ones. } ",
        "introduction": "The cosmic IR background \\citep[CIB,][]{puget1996,hauser1998} accounts for roughly half of the total extragalactic background light \\citep[EBL,][]{hauser2001,lagache2005}, i.e., half of the energy radiated by all galaxies, at all cosmic epochs, at any wavelength \\citep{dole2006}. It is therefore a crucial constraint on modes and times of galaxy formation.  Deep cosmological surveys carried out with the {\\em Infrared Space Observatory} (ISO, see Genzel \\& Cesarsky \\citeyear{genzel2000} for a summary) and {\\em Spitzer Space Telescope} \\citep[][for a review]{soifer2008} produced large samples of mid-IR sources and deep number counts \\citep{elbaz2002,papovich2004}. These surveys led to mid-IR CIB lower limits within a factor of two from the upper constraints set by TeV cosmic opacity measurements \\citep[e.g.][]{franceschini2009}.  However, at CIB peak wavelengths ($100-200$ $\\mu$m), the nature of individual galaxies building up the EBL is poorly known. Past surveys produced limited samples of distant far-IR objects \\citep[e.g.][]{frayer2009}, mainly due to the small apertures of the available instruments and the low sensitivity in the far-IR. In the 160 $\\mu$m Spitzer/MIPS band, $\\sim$7\\% of the CIB was resolved into individually detected objects \\citep{dole2004}, and it was only through stacking the 24 $\\mu$m sources that most ($60-70$\\%) of the far-IR CIB could be recovered \\citep{dole2006,bethermin2010}. Similarly, at longer wavelengths, only stacking of 24 $\\mu$m sources on BLAST maps could account for the majority of the EBL at 250, 350, and 500 $\\mu$m \\citep{marsden2009}.  With the favorable diffraction limit of the large {\\em Herschel} 3.5 m mirror (Pilbratt et al., \\citeyear{pilbratt2010}), and the high sensitivity of its {\\em Photodetector Array Camera \\& Spectrometer} (PACS, 70, 100, 160 $\\mu$m; Poglitsch et al., \\citeyear{poglitsch2010}), confusion and blending of sources are much less of a limitation. We are now able to resolve a large fraction of the CIB at its peak into individual galaxies.  The {\\em PACS Evolutionary Probe} (PEP) extragalactic survey samples 4 different tiers: from the wide and shallow COSMOS field, through medium size areas like the Lockman Hole, all the way down to the 160 and 100 $\\mu$m confusion limit in the pencil-beam, very deep observations in GOODS-N and GOODS-S, and even beyond by exploiting gravitational lensing in low-redshift galaxy clusters (e.g. Abell 2218, Altieri et al. \\citeyear{altieri2010}).    \\begin{figure*}[!ht] \\centering \\rotatebox{-90}{ \\includegraphics[height=0.425\\textwidth]{14610f1a.eps} } \\rotatebox{-90}{ \\includegraphics[height=0.425\\textwidth]{14610f1b.eps} } \\caption{Number counts at 100 ({\\em left}) and 160 $\\mu$m ({\\em right}), normalized to the Euclidean slope. Filled/open symbols belong to flux bins above/below the 80\\% completeness limit. Models belong to \\citet{lagache2004}, \\citet{franceschini2009}, \\citet{rowanrobinson2009}, \\citet{leborgne2009}, \\citet{valiante2009}, \\citet{lacey2009}, Gruppioni \\& Pozzi (in prep.). Shaded areas represent ISO and Spitzer data \\citep{rodighiero2004,heraudeau2004,bethermin2010}. {\\em Inset}: Collection of PACS number counts, including this work (red, limited to relative errors $\\le$15\\% for clarity sake), PEP Abell 2218 (grey, Altieri et al. \\citeyear{altieri2010}) and HerMES-PACS (light-blue, Aussel et al. \\citeyear{aussel2010}). The solid lines in the insets mark the trends expected for a non-evolving population of galaxies.} \\label{fig:pep_counts} \\end{figure*}  Here we exploit the {\\em science demonstration phase} (SDP) observations of the GOODS-N field, complemented with the COSMOS and Lockman Hole wider and shallower layers of the survey, to build far-IR galaxy number counts and derive the fraction of CIB resolved by PEP. We then take advantage of the extensive multi-wavelength coverage of the GOODS-N field, with the aim of identifying the CIB contributors and the epoch when the bulk of the Universe IR energy budget was emitted. ",
        "conclusions": "\\label{sect:discussion}   \\begin{figure*}[!ht] \\centering \\includegraphics[height=0.4\\textwidth]{14610f3a.eps} \\includegraphics[height=0.4\\textwidth]{14610f3b.eps} \\caption{Slicing the PEP GOODS-N population into redshift bins. Each PACS source has been associated to an optical/mid-IR counterpart as described in Appendix \\ref{sect:data}. See Fig. \\ref{fig:pep_counts} and text for details on the models shown.} \\label{fig:pep_counts_redshift} \\end{figure*}  In the attempt to reproduce the observed ISO and Spitzer number counts, several authors built ``backward'' evolutionary models, including luminosity and/or density evolution, as well as different galaxy populations. In Fig. \\ref{fig:pep_counts}, we overlay recent models onto the observed PACS counts. We include in this collection also the \\citet{lacey2009} $\\Lambda$-CDM semi-analytical model (SAM), complemented with radiative transfer dust reprocessing.  The most successful models are the \\citet{valiante2009}, including luminosity-dependent distribution functions for the galaxy IR SEDs and their AGN contribution, and \\citet{rowanrobinson2009}, employing analytic evolutionary functions without discontinuities and 4 galaxy populations.  Gruppioni and \\citet{lagache2004} overestimate the amplitude of the number counts peak in both bands, while \\citet{franceschini2009} and \\citet{leborgne2009} reproduce the counts fairly well only in one channel (100 and 160 $\\mu$m, respectively). It is worth recalling that to date most of these models have been fine-tuned to reproduce mid-IR and sub-mm statistics, while a big gap in wavelength was affecting far-IR predictions. Most include a luminosity evolution $\\propto(1+z)^{3.0-3.5}$, but the redshift limit for this slope, the details of density evolution, or the adopted galaxy zoo vary significantly from author to author.  Besides spanning a much wider range of observational data (UV, optical, near-IR luminosity functions, galaxy sizes, metallicity, etc.), the SAM approach suffers here for a limited flexibility in the choice of parameters, and significantly overestimates the bright end of PACS counts. Moreover, it cannot reproduce the peak, especially in the green band.  Thanks to the rich ancillary dataset in GOODS-N, we split the far-IR number counts into redshift bins (Fig. \\ref{fig:pep_counts_redshift}). This elaboration offers a remarkable chance to set detailed constraints on the evolution of the galaxy populations adopted in current recipes. This view highlights some new features and the main problems of the models under discussion.  First, in the lowest redshift bin, $0.0<z\\le0.5$, the differential number counts, normalized to the Euclidean slope, monotonically increase as a function of flux, resembling the trend expected for a non-evolving population of galaxies. The consistency of models with data weakens at longer wavelengths. At the bright end, above 10 mJy and 20 mJy at 160$\\mu$m, PEP catalogs include 43 and 18 objects in this redshift range. Based on template fitting (Rodighiero et al., \\citeyear{rodighiero2010}), half of this sample shows a late-spiral SED ({\\em Sd-Sdm} in the Polletta et al., \\citeyear{polletta2007}, template library) with typical luminosities above $10^{11}$ L$_\\odot$ and a peak wavelength of $\\sim$130 $\\mu$m. Nevertheless, at these luminosities, \\citet{chary2001} models, which adopt the local $L-T$ relation, predict SEDs peaking around 90 $\\mu$m restframe, thus indicating that this class is indeed colder than typical local IR-luminous sources. This is reminiscent of the cold sources found by ISO FIRBACK and serendipity surveys \\citep[e.g.][]{dennefeld2005,stickel2007}.    Significant differences between models arise at higher redshift. The bin $1.0<z\\le2.0$ shows the largest discrepancies: apparently, most of the overestimation of total number counts by some models originates here, suggesting that the adopted evolution is either too protracted in time or too steep.  We are now able for the first time to infer how much of the resolved CIB peak was emitted at different cosmic epochs. Figure \\ref{fig:pep_cib} depicts the amount of extragalactic IR background originating from the four redshift bins taken into account (see also Fig. \\ref{fig:pep_counts_redshift}). At 100 $\\mu$m, at the GOODS-N 3$\\sigma$ limit of $\\sim$3 mJy, roughly 80\\% of the resolved CIB is emitted by objects detected at $z\\le1$, equally distributed below and above $z=0.5$. The remainder belongs to more distant galaxies, with only 5\\% locked into individual objects at $z>2$. Most of the resolved 160 $\\mu$m CIB is produced at higher redshift, with the $0.5<z\\le1.0$ redshift bin dominating ($\\sim$40\\%) the budget.  At the PEP depth, the CIB resolved into individual sources is mainly due to luminous and ultraluminous IR galaxies (LIRGs $10^{11}\\le L_{IR}<10^{12}$ L$_\\odot$, ULIRGs $L_{IR}\\ge10^{12}$ L$_\\odot$). The relative fraction of the two classes varies as a function of redshift, likely due to selection effects, with LIRGs providing 80-90\\% of the CIB resolved at $0.5<z\\le1$ and ULIRGs dominating at higher redshift (60-70\\% at $1.0<z\\le2.0$ and 100\\% above).  Finally, Fig. \\ref{fig:pep_cib} also reports the CIB contribution of 24 $\\mu$m sources, as obtained through stacking (see Sect. \\ref{sect:cib}). The relative CIB fractions emitted in the four redshift bins change significantly at 100 $\\mu$m, with respect to those obtained by integrating the observed counts. The redshift distribution of the CIB is now peaking between $0.5<z\\le1.0$ and the relative contribution of $z>1$ increases by more than 10\\%. On the other hand, at 160 $\\mu$m only the relative fraction of the highest redshift bin varies by more than 5\\%, while the others remain unchanged within $1-2$\\%. Also in the case of stacking, the 160 $\\mu$m CIB relative redshift distribution is more populated at high redshift than the 100 $\\mu$m one. As expected, this indicates that PACS galaxies with redder observed colors lie on average at higher redshift than bluer ones.      \\onltab{1}{ \\begin{table*}[!ht] \\centering \\begin{tabular}{r | c c | c c | c c | c c | c c | c c} \\hline \\hline \\multicolumn{1}{c|}{$S_{center}$} & \\multicolumn{2}{c|}{GOODS-N 100$\\mu$m} & \\multicolumn{2}{c|}{LH 100$\\mu$m} & \\multicolumn{2}{c|}{COSMOS 100$\\mu$m} & \\multicolumn{2}{c|}{GOODS-N 160$\\mu$m} & \\multicolumn{2}{c|}{LH 160$\\mu$m} & \\multicolumn{2}{c}{COSMOS 160$\\mu$m}\\\\ \\hline 2.84  &  3.19e+04   &    0.208 &      --  \t &    --    & \t    --\t     &    --\t&   --\t     &    --\t&    \t--\t&\t-- &\t --\t   &\t-- \\\\ 3.57  &  4.67e+04   &    0.146 &      --  \t &    --    & \t    --\t     &    --\t&   --\t     &    --\t&    \t--\t&\t-- &\t --\t   &\t-- \\\\ 4.50  &  5.57e+04   &    0.149 &      --  \t &    --    & \t    --\t     &    --\t&   --\t     &    --\t&    \t--\t&\t-- &\t --\t   &\t-- \\\\ 5.66  &  6.37e+04   &    0.162 &      7.00e+04\t &  0.086   & \t    --\t     &    --\t&   1.17e+05 &    0.129 &    \t--\t&\t-- &\t --\t   &\t-- \\\\ 7.13  &  6.90e+04   &    0.185 &      7.22e+04\t &  0.100   & \t  7.23e+04   &    0.028\t&   1.42e+05 &\t  0.173 &    \t--\t&\t-- &\t --\t   &\t-- \\\\ 8.97  &  6.43e+04   &    0.234 &      6.58e+04\t &  0.128   & \t  8.10e+04   &    0.031\t&   1.63e+05 &\t  0.142 &    1.43e+05\t&   0.086  &\t --\t   &\t-- \\\\ 11.30  &  6.29e+04   &    0.287 &      7.24e+04\t &  0.146   & \t  7.47e+04   &    0.039\t&   1.74e+05 &\t  0.163 &    1.67e+05\t&   0.093  &\t --\t   &\t-- \\\\ 14.22  &  6.81e+04   &    0.333 &      8.13e+04\t &  0.167   & \t  6.81e+04   &    0.050\t&   1.66e+05 &\t  0.200 &    1.82e+05\t&   0.106  &\t --\t   &\t-- \\\\ 17.91  &  1.11e+05   &    0.302 &      7.00e+04\t &  0.217   & \t  7.33e+04   &    0.057\t&   1.77e+05 &\t  0.237 &    1.70e+05\t&   0.132  &\t 2.03e+05  &\t0.033\\\\ 22.54  &  8.13e+04   &    0.443 &      6.66e+04\t &  0.269   & \t  7.15e+04   &    0.071\t&   1.97e+05 &\t  0.266 &    1.53e+05\t&   0.171  &\t 2.07e+05  &\t0.039\\\\ 28.38  &  1.11e+05   &    0.438 &      6.70e+04\t &  0.317   & \t  6.70e+04   &    0.087\t&   1.83e+05 &\t  0.334 &    1.85e+05\t&   0.183  &\t 1.95e+05  &\t0.049\\\\ 35.73  &  6.63e+04   &    0.701 &      4.01e+04\t &  0.529   & \t  5.74e+04   &    0.116\t&   1.88e+05 &\t  0.394 &    1.66e+05\t&   0.238  &\t 1.80e+05  &\t0.062\\\\ 44.98  &  6.68e+04   &    0.814 &      5.97e+04\t &  0.484   & \t  5.91e+04   &    0.135\t&   1.49e+05 &\t  0.527 &    1.42e+05\t&   0.312  &\t 1.48e+05  &\t0.082\\\\ 56.62  &  --\t  &   --     &    $\\le$1.81e+04\t &  1.000   & \t  5.06e+04   &    0.175\t&   1.44e+05 &\t  0.641 &    1.14e+05\t&   0.415  &\t 1.41e+05  &\t0.100\\\\ 71.29  &  --\t  &   --     &\t  $\\le$2.40e+04\t &  1.000   & \t  5.64e+04   &    0.198\t&   1.01e+05 &\t  0.883 &    7.01e+04\t&   0.658  &\t 1.16e+05  &\t0.135\\\\ 89.74  &  --\t  &   --     &\t    8.46e+04\t &  0.655   & \t  3.89e+04   &    0.289\t&   --       &\t  --    &    6.37e+04\t&   0.828  &\t 9.19e+04  &\t0.185\\\\ 112.98  &  --\t  &   --     &\t\t--\t &   --\t    & \t  2.51e+04   &    0.453\t&   --\t     &    --\t&    8.19e+04\t&   0.833  &\t 7.70e+04  &\t0.239\\\\ 142.23  &  --\t  &   --     &\t\t--\t &   --\t    & \t  3.81e+04   &    0.412\t&   --\t     &    --\t&       --\t&       -- &\t 8.38e+04  &\t0.271\\\\ 179.06  &  --\t  &   --     &\t\t--\t &   --\t    & \t    --\t     &    --\t&   --\t     &    --\t&       --\t&       -- &\t 6.57e+04  &\t0.366\\\\ \\hline \\end{tabular} \\tablefoot{Fluxes are provided in $[$mJy$]$. For each field/wavelength, we list counts in units of $[$deg$^{-2}$ mJy$^{1.5}]$. Errors are given as relative fractions, and include both Poisson statistics and propagation of photometric uncertainties.} \\caption{PEP number counts, normalized to the Euclidean slope.} \\label{tab:pep_counts} \\end{table*} }"
    },
    "1005/1005.0149_arXiv.txt": {
        "abstract": "{} { An optical search is carried out for supernova remnants (SNRs) in the Sc type nearby spiral galaxy M74, using ground-based observations at TUBITAK National Observatory (TUG, Antalya/Turkey) and Special Astrophysics Observatory (SAO, Russia). Observations are supplemented by the spectral analysis of archived X-ray data from $XMM$-$Newton$\\ and $Chandra$.} { A survey of M74 covering $\\sim$ 9 arcmin$^{2}$ with [S II], H$\\alpha$, and their continuum filters. Interference filter images of M74 are obtained with 1.5 m Russian Turkish Telescope (RTT150) at TUG and  spectral data are taken with the 6 m Bolsoi Azimuthal Telescope (BTA) at SAO. The emission nebulae with continuum-subtracted line ratio values of [S II]$\\lambda$$\\lambda$6716,6731 /H$\\alpha$ $\\geq$ 0.4 are identified as SNRs. A follow-up spectroscopy is obtained to confirm optical SNR identifications.} { We have identified nine new SNR candidates in  M74 with [S II]/H$\\alpha$ $\\geq$ 0.4 as the basic criterion. We obtain [S II]/H$\\alpha$ ratio in the range from 0.40 to 0.91 and H$\\alpha$ intensities from 2.8 $\\times$ $10^{-15}$ erg cm$^{-2}$ s$^{-1}$ to 1.7 $\\times$ $10^{-14}$ erg cm$^{-2}$ s$^{-1}$. We also present spectral follow-up observations of these  SNR candidates. However, we are able to spectrally confirm only three  of them (SNR2, SNR3, and SNR5). The lack of confirmation for the rest might be due to the contamination by the nearby H II emission regions as well as due to the inaccurate positioning of the long slit on these objects. In addition, we search the $XMM$-$Newton$\\ and $Chandra$ Observatory archival data for the X-ray counterparts to the optically identified candidates. We find positional coincidence with only three SNR candidates, SNR1, SNR2, and SNR8. The spectrum of SNR2 yields a shock temperature of 10.8 keV with an ionization timescale of 1.6 $\\times$ 10$^{10}$ s cm$^{-3}$ indicating a relatively young remnant in an early Sedov phase which is not supported by our optical wavelength analysis. Given the high luminosity of 10$^{39}$ erg s$^{-1}$ and the characteristics of the X-ray spectrum, we favor an  Ultra Luminous X-ray Source interpretation for this source associated with an SNR. We calculate  an X-ray flux upper limit of 9.0 $\\times$ $10^{-15}$ erg cm$^{-2}$ s$^{-1}$ for the rest of the SNRs including spectroscopically identified SNR3 and SNR5.} {} ",
        "introduction": "M74 (NGC 628) is an Sc type spiral galaxy with a $\\sim$ 6$^{\\circ}$ inclination angle at an assumed distance of 7.3 Mpc (Sharina et al.1996, Soria et al. 2004). It is the brightest member of the small M74 group in the Pisces constellation.  M74 has spiral arms well formed with bright blue star clusters and dark cosmic dust lanes. In the early 2000, two supernovae (SN 2002ap and SN 2003gd) were identified in M74 and studied extensively. The supernova event SN 2002ap is one of the four Type Ic SNe observed in X-rays (Van Dyk et al. 2003). SN 2003gd, which is found to be a nearby Type II-P (plateau) event has a confirmed red giant progenitor (Hendry et al. 2005).  Leievre \\& Roy (2000) have recovered well over 100 small and isolated H II regions and measured their fluxes in the extreme outer disk of M74 with R$>$R$_{25}$ making use of deep H$\\alpha$ narrowband imaging. They show that massive star formation rate measured by the azimuthally averaged H$\\alpha$ surface brightness decreases monotonically from the center out to R$\\sim$20 kpc beyond which it drops rapidly.  Elmegreen et al. (2006) have examined the size and luminosity distributions of the star forming regions in the galaxy with the Hubble Space Telescope. Their results suggest clumping of stars having a scale free nature in the galaxy disk. Fathi et al. (2007) have studied the internal kinematics of M74 over its entire face using Fabry-Perot interferometry with a good angular resolution, and confirmed the presence of an inner rapidly rotating disc-like component caused by the slow secular evolution of the large-scale spiral arms together with the oval structure. They also detect over 300 H II regions in the galaxy with calibrated luminosities and diameters up to about 300 pc using continuum subtracted narrow band images in H$\\alpha$.  M74 is observed in the X-rays using the $Chandra$ and $XMM$-$Newton$\\ data by Soria et al. (2004). 74 discrete X-ray sources are detected in the combined $Chandra$ observations to a detection limit of $\\approx$ 6 $\\times$ 10$^{36}$ erg s$^{-1}$. They estimate 15$\\%$ of the M74 discrete sources to have soft colors. They also calculate thermal SNRs with luminosities $\\approx$ 2 $\\times$ 10$^{37}$ erg s$^{-1}$. Two bright  X-ray sources (XMMU J013636.5+155036, CXOU J13651.1+154547) are detected by $XMM$-$Newton$\\ in M74 (Soria et al. 2004), while only the latter is detected by Chandra (Krauss et al. 2003). It is established that both of the sources are variable. It is also argued that they can be Ultra Luminous X-ray sources (ULX), however their true nature remains unclear.   Radio observations of M74 have been carried out during the Arecibo Galaxy Environment Survey (AGES). This survey is designed to investigate M74  group environment in 21 cm with a better sensitivity and higher spatial and velocity resolution than previous observations (Auld et al. 2006). They obtain the spatial distribution of HI for M74 and other selected galaxies.  In the UV regime, the surface brightness and color profiles for M74  are calculated by Cornett et al. (1994). They show that M74 disk has sustained significant star formation over the last five hundred million years.  In this paper, we present our results on the search for new supernova remnant (SNR) candidates in M74 using the observations by RTT150 and 6 m-BTA telescopes. We also make use of archived $Chandra$ data to look for their X-ray counterparts. SNR studies are important in the theories of interstellar medium (ISM), and star formation since supernovae inject large amounts of matter, and energy into the ISM. However, in spite of the quite large number of Galactic SNRs their observations are impeded by several limitations such as the uncertainty in distances to individual objects and high extinctions along the line of sight in several regions of the Galactic plane. The limitations and uncertainties are inherently much less in extragalactic samples. Assuming that all SNRs are of the same distance to us for a given galaxy, we can easily compare  their observed properties. Relative positions of such SNRs are determined with more precision. Once we know the positions of SNRs, their distances with respect to H II regions and spiral arms can easily be calculated (Matonick \\& Fesen 1997; Blair \\& Long 1997, 2004).  SNRs are identified in a number of nearby spiral galaxies using optical observations (e.g. D$^\\prime$Odorico, Dopita \\& Benveuti 1980; Braun \\& Walterbos 1993; Magnier et al. 1995; Matonick \\& Fesen 1997; Matonick et al. 1997; Gordon et al. 1998,1999; Blair \\& Long 1997, 2004; Sonbas et al. 2009), X-ray observations (Pence et al. 2001; Ghavamian et al. 2005), and radio observations (Lacey et al. 1997; Lacey \\& Duric 2001; Hyman et al. 2001). Multiwavelength surveys of SNRs are also carried out by Pannuti et al. (2000, 2002, 2007).  In our work, we use a well known and accepted criterion (S II/H$\\alpha$ $\\geq$ 0.4) proposed by Mathewson \\& Clarke (1973) to differentiate SNRs from typical H II regions. Rationale for this lies in the fact that in a typical H II region, sulfur is usually expected to be in the form of $S^{++}$ due to strong photoionizing fluxes from central hot stars,  making the ratio [S II]/H$\\alpha$, typically, in the range $\\sim$ 0.1 - 0.3. Shock waves produced by  SN explosion propagate through the surrounding medium. The matter cools sufficiently behind these waves and variety of ionization states occur including S$^{+}$. This might be the reason for the increased [S II]/H$\\alpha$ ratio observed in SNRs.  It follows that almost all discrete emission nebulae satisfying the above criterion are expected to be shock-heated.  The organization of the paper is as follows; In section 2, we discuss our imaging and spectroscopic observations and related data reduction. Search results in the optical band, the identification of SNRs, and search for their X-ray counterparts are presented in section 3. Finally, conclusions and discussions of our results are provided in section 4. ",
        "conclusions": "We conducted a survey of SNRs using optical imaging and spectroscopic measurements in the nearby spiral galaxy M74. In this survey, we used blinking between continuum-subtracted H$\\alpha$ and continuum-subtracted [S II] images to deduce SNR candidates for a given ratio between the lines.  The SNRs were confirmed by additional spectroscopic observations and finally, comparing the optically detected SNRs with archived Chandra observations. Nine SNRs were identified in M74 by our method. Three of these have shown positional coincidences with X-ray sources. SNR2 is our best candidate to be an SNR using optical wavelength analysis and the X-ray source is consistent with a ULX interpretation, making it one of the few associations between a SNR and a ULX. We calculated an X-ray flux upper limit of 9.0 $\\times$ 10$^{-15}$ erg cm$^{-2}$ s$^{-1}$ and a luminosity upper limit of 5.0 $\\times$ 10$^{37}$ erg s$^{-1}$ for our new SNR candidates excluding SNR1, SNR2 and SNR8. Our spectroscopically detected SNRs have shock velocities $<$ 85 kms$^{-1}$  indicating that they fall into a low shock velocity range implying very old remnants from which  little X-ray emission is expected.  According to the SN rate calculation of Matonick \\& Fesen (1997) in such galaxies, about half of all SNe are of Type Ib/c or Type II, produced by massive stars. In turn, only half of these SNe are located in regions with enough ambient density to produce a detectable SNR. This means that, only about a quarter of all SN events may leave easily detectable optical remnants. In this respect four times more SN events should have exploded in M74 and we detected nine SNRs. With this reasoning, we would expect to see $\\sim$ 40 SNRs. If the optically observable lifetime of an SNR is assumed to be $\\sim$ 20,000 years (Braun, Goss, \\& Lyne 1989), a crude SN occurrence rate for M74 can be estimated giving a value of about 1 per $\\sim$ 500 years. In the light of observation of two recent SNe from M74, our result can be reconciled if we accept the fact that more SNe occur at places that are difficult to reach optically, (i.e. in gas-rich regions of the galaxy).  \\begin {acknowledgements} We thank the TUBITAK National Observatory (TUG) and the Special Astrophysical observatory (SAO) for their support for observing times and equipments. \\end {acknowledgements}"
    },
    "1005/1005.2466_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1}  In 1998, two independent supernovae (SN) observation groups found that our universe is undergoing an accelerated expansion at the present stage, through the observations of distant type Ia supernovae~\\cite{SN}. This implies that there exists a mysterious component, dark energy, which has large enough negative pressure, responsible for the cosmic acceleration. Many other astronomical observations, such as surveys of the large scale structure (LSS)~\\cite{LSS} and measurements of the cosmic microwave background (CMB) anisotropy~\\cite{CMB}, also firmly indicate that dark energy is the dominant component in the present-day universe. It is commonly believed that exploring the nature of dark energy is one of the focuses in the realm of both cosmology and theoretical physics today.  The most obvious candidate for dark energy is the famous Einstein's cosmological constant $\\lambda$ which has the equation of state $w=-1$. However, as is well known, the cosmological constant is plagued with the ``fine-tuning'' and ``cosmic coincidence'' problems \\cite{problems}. Another promising candidate for dark energy is the dynamical scalar field, a slowly varying, spatially homogeneous component. An example of scalar-field dark energy is the so-called quintessence\\cite{quintessence}, a scalar field $\\phi$ slowly evolving down its potential $V(\\phi)$. Provided that the evolution of the field is slow enough, the kinetic energy density is less than the potential energy density, giving rise to the negative pressure responsible to the cosmic acceleration. So far, in order to alleviate the cosmological-constant problems and explain the accelerated expansion, a wide variety of scalar-field dark energy models have been proposed. Besides quintessence, these also include phantom, $k$-essence, tachyon, ghost condensate and quintom amongst many. However, we should note that the mainstream viewpoint regards the scalar-field dark energy models as a low-energy effective description of the underlying theory of dark energy.  It is generally believed by theorists that we cannot entirely understand the nature of dark energy before a complete theory of quantum gravity is established. However, although we are lacking a quantum gravity theory today, we still can make some efforts to explore the nature of dark energy according to some principles of quantum gravity. The holographic dark energy model \\cite{Li:2004rb} is just an appropriate example, which is constructed in light of the holographic principle of quantum gravity theory. That is to say, the holographic dark energy model possesses some significant features of an underlying theory of dark energy. More recently, a new model consistent with the holographic principle, the agegraphic dark energy model, has been proposed in ~\\cite{Cai:2007us}, which takes into account the uncertainty relation of quantum mechanics together with the gravitational effect in general relativity.    While, by far, a complete theory of dark energy has not been established presently, we can, however, speculate on the underlying theory of dark energy by taking some principles of quantum gravity into account. The agegraphic dark energy model is no doubt a tentative in this way. Now, we are interested in that if we assume the holographic/agegraphic vacuum energy scenario as the underlying theory of dark energy, how the low-energy effective scalar-field model can be used to describe it. In this direction, some work has been done, see, e.g., \\cite{holoscalar,Zhang:2009un,Zhang:2008mb,Cui:2009ns}. The agegraphic versions of scalar-field models, such as quintessence and tachyon, have been constructed \\cite{Zhang:2008mb,Cui:2009ns}. In this paper, we focus on the canonical scalar-field description of the agegraphic dark energy, namely, the ``agegraphic quintessence'' \\cite{Zhang:2008mb}.  In recent years, cosmological-constant/dark-energy problem has been studied by string theorists within the string framework. It is generally considered that string theory is the most promising consistent theory of quantum gravity. Based on the KKLT mechanism \\cite{Kachru:2003aw}, a vast number of metastable de Sitter vacua have been constructed through the flux compactification on a Calabi-Yau manifold. These string vacua can be described by the low-energy effective theories. Furthermore, it is realized that the vast series of semiclassically consistent field theories are actually inconsistent. These inconsistent effective field theories are believed to locate in the so-called ``swampland'' \\cite{Vafa:2005ui}. The self-consistent landscape is surrounded by the swampland. Vafa has proposed some criterion to the consistent effective field theories \\cite{Vafa:2005ui}. Moreover, it was conjectured by Arkani-Hamed et al. \\cite{weakgrav} that the gravity is the weakest force, which helps rule out those effective field theories in the swampland. Arkani-Hamed et al. pointed out \\cite{weakgrav} that when considering the quantum gravity, the gravity and other gauge forces should not be treated separately. For example, in four dimensions a new intrinsic UV cutoff for the U(1) gauge theory with single scalar field, $\\Lambda=g M_{\\rm{p}}$, is suggested, where $g$ is the gauge coupling \\cite{weakgrav}. In \\cite{hls}, the weak gravity conjecture together with the requirement that the IR cutoff should be smaller than the UV cutoff leads to an upper bound for the cosmological constant. In addition, for the inflationary cosmology, the application of the weak gravity conjecture shows that the chaotic inflation model is in the swampland \\cite{Huangchaotic}. This conjecture even implies that the eternal inflation may not be achieved \\cite{HLW}. Furthermore, Huang conjectured \\cite{Huang:2007qz} that the variation of the inflaton should be smaller than the Planck scale $M_{\\rm{p}}$, and this can make stringent constraint on the spectral index.  More recently, the weak gravity conjecture has been applied to the dark-energy problem. It is suggested that the variation of the quintessence field value $\\phi $ should be less than $M_{\\rm{p}}$. This criterion may give important theoretical constraints on the equation-of-state parameter of quintessence models, and some of these constraints are even stringent than those of the present experiments \\cite{Huang:2007mv}. The criterion $|\\Delta\\phi(z)|/M_{\\rm{p}}\\leq 1$ has also been used to put theoretical constraints on other canonical scalar-field dark energy models; see, e.g., \\cite{Ma:2007av,Wu:2007pq,Chen:2008gi}. In this paper we shall investigate the possible theoretical limits on the parameter $n$ of the agegraphic quintessence from the weak gravity conjecture.  In the next section, we will briefly review the new agegraphic dark energy model proposed in~\\cite{Wei:2007ty}. In Sec.~\\ref{sec3}, we will give the possible theoretical limits on the parameter $n$ of the agegraphic quintessence model from the weak gravity conjecture. Conclusion will be given in Sec.~\\ref{sec4}.  \\setcounter{equation}{0} ",
        "conclusions": "\\label{sec4}  To summarize, in this paper we have investigated the theoretical limits on the parameter $n$ of the agegraphic quintessence model by considering that the variation of the quintessence scalar field $\\phi$ should be less than the Planck mass $M_{\\rm{p}}$. The agegraphic dark energy can mimic the behavior of a quintessence scalar-field dark energy, so the quintessence model can be used to effectively describe the agegraphic dark energy. In this paper, we have tested the single-field and multi-field agegraphic quintessence models by using the weak gravity conjecture. We believe that the low-energy effective field theory is not applicable in the trans-Planckian field space.  We have shown that for both single-field and multi-field agegraphic quintessence models the weak gravity conjecture leads to the same theoretical limit, $n\\lesssim 2.5$, which is inconsistent with the current observational constraint $2.637<n<2.983$ (95.4\\% CL). The requirement that the variation of the field should be less than the Planck scale from weak gravity conjecture may arise from the consistent theory of quantum gravity, so in this sense the theoretical result obtained in this paper can, to some extent, be viewed as the prediction of quantum gravity. The tension between theoretical and observational limits implies that the agegraphic dark energy could not be described by a consistent low-energy effective scalar field theory. Of course, other possible reasons for the tension still exist, for example, perhaps the current observational data cannot provide precise constraint on the agegraphic dark energy model, and the future accurate data might change the constraint result to be consistent with the theoretical limit from weak gravity conjecture."
    },
    "1005/1005.5300_arXiv.txt": {
        "abstract": "We report the discovery of \\hatcurb{}, a transiting extrasolar planet in the `period valley', a relatively sparsely-populated period regime of the known extrasolar planets. The host star, \\hatcurCCgsc\\,, is a \\hatcurISOspec\\, dwarf with $V=$\\hatcurCCtassmvshort\\,. It has a mass of \\hatcurISOm\\,\\msun, radius of \\hatcurISOr\\,\\rsun, effective temperature \\hatcurSMEteff\\,K, and metallicity $\\feh =\\hatcurSMEzfeh$. The planetary companion orbits the star with a period $P=\\hatcurLCP\\,d$, transit epoch $T_c =\\hatcurLCT$ (BJD), and transit duration \\hatcurLCdur\\,d. It has a mass of \\hatcurPPmlong\\,\\mjup, and radius of \\hatcurPPrlong\\,\\rjup\\ yielding a mean density of \\hatcurPPrho\\,\\gcmc. At an age of \\hatcurISOage\\,Gyr, the planet is H/He-dominated and theoretical models require about $2$\\% ($10$\\mearth) worth of heavy elements to reproduce its measured radius. With an estimated equilibrium temperature of $\\sim820$~K during transit, and $\\sim1000$~K at occultation, \\hatcurb{} is a potential candidate to study moderately cool planetary atmospheres by transmission and occultation spectroscopy. ",
        "introduction": "\\label{sec:introduction}  Ground-based transit surveys have been very successful at discovering short-period ($P<5$~days) transiting extrasolar planets (TEPs) since 2006. Because of the short duration of transit events (relative to the orbital period), candidates are more likely to be identified by higher duty cycle observations. The situation is especially demanding toward longer orbital periods, where the observational constraint couples with the considerably lower incidence rates of preferred orbital plane inclinations. The only way we can increase the probability of discovering long-period TEPs from the ground is to increase the duty cycle by establishing telescope networks spread in geographical longitude. Surveys such as WATTS/TrES and HATNet are designed with this goal in mind. For example, in the particular case of WATTS/TrES \\citep{oetiker:2010}, for a typical observing season, the detection probability of a 10~day period TEP from a single observing site is 1--5\\%. This increases to 10\\% by using two sites of the network. It is clear that discovering long-period TEPs with ground-based surveys involves a large degree of serendipity even in the case of multisite observations.  \\begin{figure}[h] \\plotone{\\hatcurhtr-pdistr.ps} \\caption{ Period distributions of the extrasolar planets discovered by radial velocity (RADV) method and by photometric transits (TEP -- a small fraction of TEPs was discovered by the radial velocity method and then later confirmed as TEPs). No TEPs are included in the distribution of the RADV planets. \\hatcurb\\, is situated at a relatively sparsely populated region in both diagrams. \\label{fig:pdistr}} \\end{figure}  In large part due to observational selection effects, from the more than $70$ TEPs\\footnote{See http://exoplanet.eu/, as of May 10, 2010} known today, there are only 6 planets with orbital periods longer than 6~days. Three of them \\citep[CoRoT-4b, 6b and 9b, see][]{fridlund:2010, deeg:2010} were discovered by the CoRoT mission; two, HD~17156b \\citep{barbieri:2007}, and HD~80606b \\citetext{\\citealp{moutou:2009}, \\citealp{fossey:2009}, \\citealp{garcia:2009} and \\citealp{hebrard:2010}} came from the photometric followup observations of already known planets from radial velocity (RV) surveys. The only long-period planet originally discovered by a ground-based photometric survey is WASP-8b \\citep{queloz:2010}, with a period of $8.16$~days. The period distribution of TEPs, displayed in the lower panel of \\reffigl{pdistr}, shows the very sparse population of TEPs above $5$~days ($\\log P > 0.7$). Interestingly, there are also relatively few radial velocity (RV) planets discovered in the period interval of $10$--$100$~days (top panel). The depletion of planets in this regime can be explained by the stellar-mass-dependent lifetime of the protoplanetary disk and its effect on migration \\citetext{\\citealp{burkert:2007} and \\citealp{currie:2009}}.  The 10.86~day period planet \\hatcurb{} was discovered by the Hungarian-made Automated Telescope Network \\citep[HATNet;][]{bakos:2004}. Operational since 2003, it has covered approximately 11\\% of the Northern sky, searching for TEPs around bright stars ($8\\lesssim I \\lesssim 12.5$\\,mag). HATNet employs six wide field instruments: four at the Fred Lawrence Whipple Observatory (FLWO) in Arizona, and two on the roof of the Submillimeter Array hangar (SMA) of SAO in Hawaii.  Since 2006, HATNet has announced and published 14 TEPs. This work describes the fifteenth such discovery, the first TEP with a period above $10$~days discovered by a ground-based photometric survey. ",
        "conclusions": "\\label{sec:discussion}  The transiting extrasolar planet (TEP) \\hatcurb{} reported in this paper is among the few with orbital periods greater than $8$~days.\\footnote{Currently, there are no TEPs between $5.6$ and $8$~days, and, together with \\hatcurb{}, there are only four TEPs with periods greater than $10$~days. The periods are more densely distributed below $5$~days.} We summarize the properties of the currently known long-period TEPs in \\reftabl{ltep}.  \\ifthenelse{\\boolean{emulateapj}}{ \\begin{deluxetable*}{lccrrcccrcccccr} }{ \\begin{deluxetable}{lccrrcccrcccccr} \\rotate } \\tabletypesize{\\scriptsize}  \\tablewidth{0pc} \\tablecaption{Long-period TEPs\\tablenotemark{a} \\label{tab:ltep}} \\tablehead{ \\colhead{Name} & \\colhead{$\\mpl$} & \\colhead{$\\rpl$} & \\colhead{$M_{\\rm core}$\\tablenotemark{b}} & \\colhead{$T_{\\rm eq}$\\tablenotemark{c}} & \\colhead{Period} & \\colhead{a} & \\colhead{e} & \\colhead{$T_{14}$} & \\colhead{Star(\\teff)} & \\colhead{[Fe/H]} & \\colhead{Age\\tablenotemark{e}} & \\colhead{$\\mstar$} & \\colhead{$\\rstar$} & \\colhead{V} \t\\\\ \\colhead{} & \\colhead{[$\\mjup$]} & \\colhead{[$\\rjup$]} & \\colhead{[\\mearth]} & \\colhead{[K]} & \\colhead{[d]} & \\colhead{[AU]} & \\colhead{} & \\colhead{[h]} & \\colhead{[K]} & \\colhead{[dex]} & \\colhead{[Gyr]} & \\colhead{[$\\msun$]} & \\colhead{[$\\rsun$]} & \\colhead{[mag]} } \\startdata WASP-8b      & 2.25    & 1.05  & 52  &  940 & 8.15872   & 0.080 & 0.31  & 3.5 & G6 (5600)  & $+0.17$ & 4.0 & 1.03 & 0.95      &  9.9 \\\\ CoRoT-6b     & 2.96    & 1.17  & 0 & 1020 & 8.88659   & 0.086 & 0.10\\tablenotemark{d}  & 3.8 & F5 (6090)  & $-0.20$ & 3.0 & 1.06 & 1.03      & 13.9 \\\\ CoRoT-4b     & 0.72    & 1.19  & 0 & 1070 & 9.20205   & 0.090 & 0.00  & 4.4 & F0 (6190)  & $+0.00$ & 1.0 & 1.10 & 1.15      & 13.7 \\\\ HAT-P-15b    & \\hatcurPPmshort & \\hatcurPPrshort & $10$ & 910 & \\hatcurLCPshort & $0.096$ & $0.19$ & $5.5$ & G5 (5568) & \\hatcurSMEiizfehshort & \\hatcurISOageshort & \\hatcurISOmshort & \\hatcurISOrshort & 12.2 \\\\ HD~17156b    &\t 3.21  & 1.02  & 136  & 880  &\t21.2169   & 0.162 & 0.68  &  3.2 & G0 (6079)  & $+0.24$ & 3.1  & 1.24 & 1.45      &  8.2 \\\\ CoRoT-9b     &\t 0.84  & 1.05  & 18   & 410  &\t95.2738   & 0.407 & 0.11  &  8.1 & G3 (5625)  & $-0.01$ & 3.0  & 0.99 & 0.94      & 13.7 \\\\ HD~80606b    &\t 4.08  & 0.98  & 195  & 400  &\t111.436   & 0.455 & 0.93  & 11.9 & G5 (5574)  & $+0.43$ & 7.6  & 1.01 & 1.01      & 8.9 \\\\ [-1.5ex] \\enddata \\tablenotetext{a}{All parameters are from {\\em http://exoplanet.eu/}, except for those of WASP-8b \\cite[see][]{queloz:2010} and HD~80606b \\citetext{see \\citealp{moutou:2009} (\\teff) and \\citealp{hebrard:2010}}.} \\tablenotetext{b}{Based on the models of \\cite{fortney:2007}; zero core mass indicates that the planet has larger radius than predicted by the models at zero core mass.} \\tablenotetext{c}{Computed at a distance given by the semi-major axis.} \\tablenotetext{d}{Upper limit.} \\tablenotetext{e}{Ages have been taken either from `exoplanet.eu' or from currently published papers. For CoRoT-6b we adopted an intermediate value from the range of $2.5$--$4.0$~Gyr listed by \\cite{fridlund:2010}. For CoRoT-9b \\cite{deeg:2010} give a wide range of $0.2$--$8$~Gyr, and we adopted a plausible intermediate-to-low age.} \\ifthenelse{\\boolean{emulateapj}}{ \\end{deluxetable*} }{ \\end{deluxetable} }   \\hatcurb{} is the longest period TEP discovered by ground-based photometric surveys. With a duration of $5.5$~hours, the transit of \\hatcurb{} is at the limit in length that can be pursued from a single site ground-based observation (the only planet with a longer period than \\hatcurb{} that can still be followed up from a single site is HD~17156b). The possibility of performing single-site full followup observations is a valuable property, since systematics are more difficult to handle when low signal-to-noise ratio data from different sites are combined. The orbit is eccentric with a high significance. Except for CoRoT-4b (and perhaps for CoRoT-6b), all of the 7 long-period planets have this property. Most of the planet masses are greater than $2\\mjup$. The host star masses are lower than $\\sim1.2$\\msun\\,, in agreement with the upper mass limit of $\\sim1.4$\\msun\\, of the extrasolar planets situated in the `period valley' between $\\sim0.1$~AU and $\\sim1$~AU \\citep{bowler:2010}.  The relative depletion of planets in the $0.1$--$1$~AU region was recognized a few years ago when the number of extrasolar planets grew to the level that made more reliable statistical investigations possible \\citep{johnson:2007}. All host stars with semi-major axes smaller than $\\sim1$~AU have masses lower than $\\sim1.4$\\msun\\,. Above $\\sim1$~AU there are several stars with $M\\sim0.5$--$2.5$\\msun\\, \\citep{bowler:2010}. The distribution of the periods (or semi-major axes) of all known extrasolar planets is bimodal with a high significance (see upper panel of \\reffigl{pdistr}). Although the sample of multiple systems is not very extensive, it is interesting to note that the distribution becomes more uniform when considering only these systems \\cite{wright:2009}. Although some aspects of these properties can be explained by engulfment of planets in evolved giants \\citep{villaver:2009}, early stage of planet formation should play more important role. According to the models considering the formation phase \\citetext{\\citealp{burkert:2007} and \\citealp{currie:2009}, see also \\citealp{kretke:2009} for the role of magnetorotational instability in the protoplanetary disk} the bimodal distribution is attributed to the stellar-mass-dependent lifetime of the protoplanetary gas disk. For higher mass stars, the gas-depletion time scale is shorter due to accretion and increased EUV radiation, and therefore planets may stop inward migration too early and get stranded on orbits at high semi-major axes. For stellar masses $\\la1.2$\\msun\\, planets do migrate. The rate of migration is determined by the ratio of the viscous and depletion time scales.\\footnote{With a careful filtering of the currently available data we found that the valley exist at all host star masses, albeit for stars with $\\ga1.2$\\msun\\, the effect is indeed much stronger.}  We investigated the agreement between the derived planetary mass and radius and predictions from current theoretical models. We interpolated the models of \\cite{fortney:2007} to the observed planet mass of $\\hatcurPPmlong$~$\\mjup$ and semi-major axis of equivalent solar irradiance (which is the same as the observed semi-major axis in this case). The iso-core lines in the $\\log(age)$--$\\log(R_{\\rm p})$ plane are shown in \\reffigl{fortney-model}. Although the errors are large, and the derived core mass (especially at the low end) depends very sensitively on the stellar/planet parameters, it seems likely that \\hatcurb{} is a H/He-dominated gas giant planet with a core mass most likely about $10$\\mearth\\,, but which could range from zero to perhaps $50$\\mearth\\,.  \\hatcurb{} does not exhibit the ``radius anomaly'', where -- as is frequently the case for hot Jupiters -- the planetary radius is significantly larger than predicted by models such as that of \\cite{fortney:2007}. Furthermore, we find that for a hypothetical planet with the same mass and radius as those of \\hatcurb{}, the derived core mass is only a weak function of the period (semi-major axis) between $\\sim10$ and $1000$~days (it varies between $0$ and $10$\\mearth\\,). However, below $10$~days the core mass increases up to $60$\\mearth\\, (at $1$~day), and starts especially being sensitive below $4$~days. These all show that \\hatcurb{} resides in the period regime where standard planet structure models (i.e., the ones already tested on Solar System planets) are satisfactory, without resorting to additional effects (e.g., extra heat source, tidal heating, enhanced atmospheric opacities, etc., see \\citealp{miller:2009}) required by many of the hot Jupiters.  In \\reftabl{ltep} we also added the derived core masses for the other long-period TEPs. We see that two of the shorter period members still exhibit some radius anomalies. For CoRoT-4b and CoRoT-6b, tidal heating \\citep{miller:2009} is unlikely, since they both have close to circular orbits. On the other hand, considering the possible ambiguities in the observed absolute planetary radii, it is possible that future followup observation will result in downward radius corrections of $3$--$5$\\% necessary to solve the apparent anomaly for these planets. This level of radius correction is not uncommon in the followup works of increasing accuracy \\citetext{e.g., the case of HAT-P-1b -- \\citealp{winn:2007} or that of HD~80606b -- \\citealp{fossey:2009} versus \\citealp{hebrard:2010}, and especially WASP-8b, with $10$\\% radius decrease in the very recent analysis of \\citealp{queloz:2010}}.  It is also useful to compare the overall dependence of the planetary radius on the incident stellar flux. \\reffigl{flux-radius} shows this dependence for TEPs in the mass range of $(0.7$--$5.0)$\\mjup. Although the stellar fluxes for the long-period TEPs are still $2$--$3$ orders of magnitude larger than that of Jupiter, it seems that this is already enough to cause a substantial decrease in the planetary radius, and bring down the otherwise inflated values to `more normal' ones, corresponding to that of the Jupiter. We note that this happens in a quite wide mass range.  \\begin{figure}[!ht] \\plotone{\\hatcurhtr_Fortney.eps} \\caption{ Position of \\hatcurb{} on the age --- planet radius diagram. The models of \\cite{fortney:2007} have been interpolated to the derived mass and solar equivalent semi-major axis of \\hatcur{}. Open circles are the tabulated/interpolated values from \\cite{fortney:2007}, continuous lines have been obtained by linear interpolation. Lines are labeled by the core mass (relative to the total planet mass). } \\label{fig:fortney-model} \\end{figure}  \\begin{figure}[!ht] \\plotone{\\hatcurhtr_radius-fluxaver.eps} \\caption{ Incident stellar flux versus planet radius for TEPs of $0.7\\la M_{\\rm p}/M_{\\rm J}\\la 5.0$. The long-period planets of \\reftabl{ltep} are shown by whiter shade (green), the triangle at the left corner shows the position of Jupiter. } \\label{fig:flux-radius} \\end{figure}   With V=$12.2$~mag, \\hatcur{} is still bright enough to be the subject of ground- and space-based observations. At a distance of $a=\\hatcurPParel$~AU from its host star, \\hatcurb{} allows us to investigate the properties of a colder planet, not influenced by the complicated hydrodynamical and radiative effects which may apply to canonical hot Jupiters (see, e.g., the discussion of the core mass insensitivity above). At the same time, the period of $\\hatcurLCPvshort$~days is still on a relatively short time scale, allowing frequent followup observations. The feasibility of possible follow-up observations is briefly summarized below. \\begin{itemize} \\item {\\em Search for other planets.} Long-term radial velocity (RV) and transit timing monitoring is justified due to the signature of excess RV scatter that is unlikely to be due to stellar activity (see \\refsecl{hispec}). Although we have made a very thorough search both in the RV and in the photometric data for a trace of a second planet, we have not found anything convincing. Due to the relatively small value of \\hatcurRVjitter\\ms\\, of the ``jitter'', the mass of the possible hidden companion is estimated in the sub-Neptune regime (except if there is a planet in, e.g., a 2:1 resonance, hiding in the eccentric solution -- see \\citealp{anglada:2010}). We note that here, the eccentricity of the orbit of \\hatcurb{} is a less forcing argument for searching for a putative second planet, because at the semi-major axis of \\hatcurb{} the circularization time scale due to tidal dissipation may be on the order of several Gyrs \\citep{matsumura:2008}. \\item {\\em Inclination of the orbital and stellar spin axes.} The measurement of the Rossiter-McLaughlin (R-M) effect for \\hatcur{} is somewhat challenging, due to its relatively low projected rotation rate of \\vsini$=$\\hatcurSMEiivsin\\,. The expected size of the effect is $\\sim20$\\ms\\,, similar to the one currently measured for WASP-8b with the same \\vsini\\, \\citep{queloz:2010}. Three of the seven long-period planets listed in \\reftabl{ltep} have yielded R-M data sufficient to compute the projected inclination of the stellar rotational and orbital axes. HD~17156 is aligned \\citep{narita:2009}, while HD~80606b is tilted \\citep{hebrard:2010}, and WASP-8b \\citep{queloz:2010} is apparently strongly tilted, suggesting a retrograde orbit. It is clear that measuring the R-M effect of long-period planets in more evolved systems such as \\hatcur{} is very important, due to the expected relaxation of the dynamics to a final equilibrium state.  \\item {\\em Thermal emission.} With a proper choice of the infrared waveband and repeated observations, the detection of the occultation of \\hatcurb{} by ground-based instruments may be feasible. Assuming blackbody radiation, at a fixed wavelength the occultation depth is equal to $\\delta_{\\rm occ}=(R_{\\rm p}/R_{\\rm star})^2(T_{\\rm pl}/T_{\\rm star})$ \\citep{winn:2010}. Assuming $T_{\\rm eq}=T_{\\rm pl}$, at the moment of occultation (which is the periastron in the case of \\hatcurb{}) we have $T_{\\rm pl}\\sim1000$~K. It follows then that $\\delta_{\\rm occ}\\sim0.002$, which is measurable, especially if we consider that the occultation lasts for $\\sim4$~hours; short-enough to cover the event in one night and long-enough to gather sufficient number of data points (especially in the case of repeated observations) to reach the level of detection of several $\\sigma$. \\item {\\em Atmospheric absorption}. Measuring the small change in the transit depth due to the varying absorption level in different wavelengths (i.e., performing transmission spectroscopy) is even more challenging than detecting the occultation event. This is because the change in the transit depth is proportional to the scale-height $H$, which is proportional to the ratio of the local temperature to the gravity $T/g$ \\citep{winn:2010}. Since \\hatcurb{} is relatively cold and has a higher gravity than `standard' hot Jupiters, it is expected that the signal will be about $25$\\% of what is usually measured (assuming `standard' hot Jupiter $T$ and $g$ of $2000$~K and $20$\\mss\\,). Considering that the relative radius variation even for short-period targets like HD~189733 is of the order of $0.1$\\% \\citep{sing:2009}, measuring absorption features in colder planets such as \\hatcurb{} would indeed be very difficult. \\end{itemize}  With the continuing work of ground-based wide-field photometric surveys, we expect to discover further TEPs in the $10$--$30$~day period regime. Since these will be mostly bright ($V\\la13$~mag) targets with orbital periods still considered to be relatively short, a wide range of followup works will be feasible. Therefore, these planets will fill the gap between the classical hot Jupiters and the long-period, but likely much fainter planetary systems (more akin to the Solar System) to be discovered by ongoing and future space projects."
    },
    "1005/1005.3526_arXiv.txt": {
        "abstract": "Processes such as the solar wind sputtering and micrometeorite impacts can modify optical properties of surfaces of airless bodies. This explains why spectra of the main belt asteroids, exposed to these `space weathering' processes over eons, do not match the laboratory spectra of ordinary chondrite (OC) meteorites. In contrast, an important fraction of Near Earth Asteroids (NEAs), defined as Q-types in the asteroid taxonomy, display spectral attributes that are a good match to OCs. Here we study the possibility that the Q-type NEAs underwent recent encounters with the terrestrial planets and that the tidal gravity (or other effects) during these encounters exposed fresh OC material on the surface (thus giving it the Q-type spectral properties). We used numerical integrations to determine the statistics of encounters of NEAs to planets. The results were used to calculate the fraction and orbital distribution of Q-type asteroids expected in the model as a function of the space weathering timescale, $t_{\\rm sw}$ (see main text for definition), and maximum distance, $r^*$, at which planetary encounters can reset the surface. We found that $t_{\\rm sw}\\sim10^6$ yr (at 1 AU) and $r^* \\sim 5$ $R_{\\rm pl}$, where $R_{\\rm pl}$ is the planetary radius, best fit the data. Values $t_{\\rm sw}<10^5$~yr would require that $r^*>20$ $R_{\\rm pl}$, which is probably implausible because these very distant encounters should be irrelevant. Also, the fraction of Q-type NEAs would be probably much larger than the one observed if $t_{\\rm sw} > 10^7$ yr. We found that $t_{\\rm sw} \\propto q^2$, where $q$ is the perihelion distance, expected if the solar wind sputtering controls $t_{\\rm sw}$, provides a better match to the orbital distribution of Q-type NEAs than models with fixed $t_{\\rm sw}$. We also discuss how the Earth magnetosphere and radiation effects such as YORP can influence the spectral properties of NEAs. ",
        "introduction": "Measurements of the spectral properties of Near Earth Asteroids (NEAs) provide important evidence concerning the relationship between asteroids and the most common class of meteorites known as the ordinary chondrites (OCs). The tendency toward seeing OC-like spectral attributes among NEAs has been noted in multi-filter color observations (Rabinowitz 1998, Whiteley 2001), and in visible and near-infrared specrophotometric surveys (Binzel et al. 1996, 2004, 2010). In contrast, no spectral analogs of OCs have been found to date among the $\\sim$2000 surveyed main belt asteroids (MBAs), except for a case related to identified recent asteroid collisions (Moth\\'e-Diniz and Nesvorn\\'y 2008).  The lack of spectrophotometric analogs for OC meteorites in the main belt is a long-debated and fundamental problem. It is now generally accepted that processes similar to those acting on the Moon, such as solar wind sputtering and micrometeorite impacts (Gold 1955, Pieters et al. 2000, see Hapke 2001 and Chapman 2004 for reviews), can darken and redden the initially OC-like spectrum of a fresh asteroid surface, giving it the `weathered' appearance (see Chapman 1996 and Clark et al. 2001, 2002a,b for direct evidence for asteroid space weathering processes from the NEAR-Shoemaker and Galileo spacecrafts). In the following text we will refer to processes that alter optical properties of surfaces of airless bodies as the `space weathering' (SW) effects.  Since the SW processes should affect the MBAs and NEAs in roughly the same way (see, e.g., Marchi et al. 2006 for a study of the SW dependence on heliocentric distance), it may seem puzzling why a significant fraction of NEAs has an unweathered appearance (Binzel et al. 2004) while practically all spectroscopically surveyed MBAs are weathered. Several explanations have been proposed.  To simplify the discussion of different models described below, we will use the following terminology taken from the standard asteroid taxonomy (Bus and Binzel 2002, DeMeo et al. 2009). We will define three categories of asteroid spectra: (1) Q-type spectra with deep absorption bands and shallow spectral slope similar to that of most OC meteorites in the RELAB database\\footnote{http://www.planetary.brown.edu/relab/}; (2) S-type spectra with shallow absorption bands and relatively steep spectral slope similar to that of weathered OCs; and (3) Sq-type spectra as the intermediate case between S and Q. See Bus and Binzel (2002) and  DeMeo et al. (2009) for a formal definition of these categories. We will assume that asteroids with the Q-type spectra have essentially unweathered surfaces with OC-like mineralogy; the Sq- and S-type asteroids will be assumed to have moderately and strongly weathered surfaces with OC-like mineralogy.  About 20\\% of chondritic NEAs surveyed at visible and/or near-infrared wavelengths have Q-type spectra while this fraction is essentially zero among MBAs. The standard interpretation of these results has been based on the presumption that Q-type asteroids are likely to be small. Indeed, the current spectrophotometric surveys of MBAs are largely incomplete in the size range of typical NEAs with diameters less then a few km. Special emphasis has therefore been given to asteroid-size-dependent processes, such as immaturity of regoliths on small asteroids and/or the shorter collisional lifetime of smaller asteroids (e.g., Johnson and Fanale 1973, Binzel et al. 1996, 2004, Rabinowitz 1998, Whiteley 2001).  One possibility is that the observed spectral variations may be related to particle-size effects (Johnson and Fanale 1973), where the decreasing gravity results in a different size distribution of surface particles on typically smaller NEAs than on larger MBAs. However, the photometric parameters indicative of particle-size effects show little evidence of an asteroid diameter dependence, thereby giving doubt to this explanation (e.g., Clark et al. 2001, Masiero et al. 2009).  Binzel et al. (2004) hypothesized that the SW size-dependency was because the survival lifetime against catastrophic disruption decreases with decreasing size. Thus, on average, as we examine smaller and smaller objects, we should see younger and younger surfaces. Surfaces showing Q-type spectral properties should thus exist, on average, only among the smallest asteroids, which become easy spectroscopic targets only when they enter into NEA space. Large, OC-like asteroids in the main belt should have, on average, space-weathered spectral properties, explaining why they are taxonomically classified as S types. ",
        "conclusions": "Several tidal effects may disturb the surface of a NEA during a distant planetary encounter.\\footnote{We only discuss distant encounters here. It is clear that the SL9-like tidal disruption, binary formation events, or events with significant mass shedding will erase any pre-existing surface features.} For example: (1) The interior structure of a rubble-pile asteroid may find a new equilibrium by re-arranging its components. This motion can produce landslides, degrade craters, ballistically displace surface material, or even remove the original layers from the asteroid. (2) Tidal stresses applied to a fractured interior may produce seismic shakes similar to, or perhaps more effective than, those generated by impacts. Consequently, surface morphology may be modified. (3) The tidal torque may spin up an asteroid. In surface segments where the centrifugal force exceeds gravity, regolith layers will be removed by carrying away the excess angular momentum. More subtle changes can occur in other surface parts of a spun-up asteroid. (4) If the tidal force becomes comparable to the object's gravity during encounter, an asteroid with large enough internal strength and a strengthless regolith may lose its regolith layer.  These effects and their dependence on the encounter distance and speed are poorly understood. Some insights into this problem can be obtained from Richardson et al. (1998), where the authors performed numerical simulations of the effects of tidal gravity on a small asteroid with strengthless (rubble-pile) interior. In the most favorable case (slow encounter speed, fast prograde rotation), they found that significant mass shedding can occur up to $\\approx$5~$R_{\\rm pl}$. This sets a soft constraint on $r^*$. On one hand, $r^*$ can be larger than 5 $R_{\\rm pl}$ because the optically-active thin surface layer may be vulnerable to even tiniest perturbations that were not considered in the Richardson et al. model. On the other hand, when averaging over all encounter geometries and plausible asteroid spin states, the mean $r^*$ can become lower than 5 $R_{\\rm pl}$. Thus, for the lack of additional constraints on $r^*$, we will tentatively assume below, as a guideline for discussion, that $r^* \\sim 5$ $R_{\\rm pl}$.  If we set $r^*=5$ $R_{\\rm pl}$ our results described in \\S3 and \\S4 imply that $t_{\\rm sw}\\sim1$~My. At first sight, this SW timescale seems to be comparable to that obtained from comparative studies of asteroid families in the main belt and OC meteorites in the RELAB database (NJWI05, Vernazza et al. 2009). For example, Vernazza et al. (2009) proposed that the SW timescale is $\\lesssim1$ My. Their result hinges on observations of two largest members of the Datura family that formed by a catastrophic breakup $\\approx$0.5 My ago (Nesvorn\\'y et al. 2006, Vokrouhlick\\'y et al. 2009). These two objects, 1270 Datura and 90265 2003CL5, appear to be significantly (but not completely) space weathered (Moth\\'e-Diniz and Nesvorn\\'y 2008), which implies that the SW timescale should be comparable to or shorter than the Datura family's age. This poses a problem because $t_{\\rm sw} \\lesssim 0.5$~My does not fit the NEA constraint (unless $r^* > 5$ $R_{\\rm pl}$). Below we discuss possible solutions to this problem.  Observations of 2001 WY35, one of the smallest known members of the Datura family (absolute magnitude $H=17$), indicate that this object is not space weathered at all (Moth\\'e-Diniz and Nesvorn\\'y 2008). If these observations were correct, they would indicate that (at least some) km-sized asteroids may weather on timescales significantly longer than $\\approx$0.5~My. For example, small km-sized fragments ejected from asteroid breakup events may not retain/accumulate sufficient regolith layer on their surface in the immediate aftermath of the collision. The SW effects may be delayed for such objects until a particulate (SW-sensitive) surface layer develops on their surface, for example, by subsequent impact shattering of the exposed rock. Thus, the regolith formation and `gardening' can be an important part of the problem (Jedicke et al. 2004, Willman et al. 2008).  We should not forget that the two constraints on the SW timescale discussed here come from studies of two distinct population of objects that are affected by different physical processes. The asteroids in the main-belt families are born by violent collisions and spend most of their lifetime beyond 2 AU. The NEAs, on the other hand, are exposed to more extreme solar-wind and temperature environment. They are olivine-rich and may therefore be more susceptible to SW effects than an average MBA (Sasaki et al. 2001, Marchi et al. 2005). While large impacts on NEAs should be rare, bombardment of their surface by $D\\sim100$ $\\mu$m particles should be more intense than on MBAs due to the larger number density of micrometeoroids at $<$2 AU (Gr\\\"un et al. 1985). Also, distant planetary encounters of NEAs should produce more gentle effects than catastrophic collisions of MBAs, thus giving the initial surface different attributes. In summary, the {\\it proper} SW timescale that measures the progression of SW under ideal conditions (e.g., in absence of regolith gardening) may be substantially shorter than the {\\it apparent} SW timescale that arises from combination of different effects, and these effects most likely operate on different timescales in the NEA and MBA environments.  Another interesting possibility is related to the effects of the Earth magnetosphere on loose particulate material on a small asteroid's surface. The Earth magnetosphere extends to $\\approx$12 $R_{\\rm Earth}$ in the direction toward the Sun, $\\approx$15 $R_{\\rm Earth}$ in apex and antapex directions, and $\\approx$25 $R_{\\rm Earth}$ the anti-helion direction. The tail region stretches well past 200 $R_{\\rm Earth}$, and the way it ends is not well-known. The magnetic field ranges from 30-60 $\\mu$T at the Earth's surface and falls roughly as $1/r^3$ with distance $r$ toward the edge of the magnetosphere. Thus, if a 100-m-sized NEA passes at distance 10 $R_{\\rm Earth}$, a 10-$\\mu$m surface dust grain subject to the Lorentz force would levitate if previously charged to $>10^8$ e. Such charge is plausible for an asteroid surface of sufficiently high electrical resistivity. It is not clear, however, whether the Lorentz force effect can be more significant than the electrostatic levitation (Lee 1996) and/or van der Waals forces (Scheeres and Hartzell 2010).  While speculative, the effects of Earth magnetosphere could possibly allow for larger $r^*$ values than those expected for tidal gravity. This could perhaps help to resolve some of the discrepancy between different measurements of the SW timescale discussed above. For example, with $r^*=20$ $R_{\\rm Earth}$, Fig. \\ref{prob2} would imply that $t_{\\rm sw} \\approx$ 0.25 My.  The orbits of Q-type NEAs in Fig. \\ref{ics} hint on bimodal distribution with a group of 7 objects with $a\\lesssim1$ AU and largely spread inclination values, and 12 objects with $a\\gtrsim1.5$ AU and $i\\lesssim10^\\circ$. Using planetary encounters as the main agent that resets the SW clock, we were not able to fit both groups simultaneously. We found that the model with fixed $t_{\\rm sw}$ can match the low-$a$ group but it fails to fit the observed fraction of Qs with $a>1.5$ AU (e.g., Fig. \\ref{m2}). On the other hand, the model with $t_{\\rm sw} \\propto q^2$ matches the high-$a$ group (Fig. \\ref{m4}) but it fails to produce the observed large fraction of Qs with $a\\lesssim1$ AU and $q<0.5$ AU.  This is puzzling. The problem may be related to biases in the current sparse spectrophotometric data. Alternatively, we may be missing some important physical effect in the model. For example, a small irregular object can be spun up by a radiation effect known as YORP and shed mass (e.g., Walsh et al. 2009; see Bottke et al. 2006 for a recent review of YORP). This could lead to a partial or global removal of the space weathered material and exposure of fresh material on the surface. This effect can therefore be important. Unfortunately, the timescale on which the surface of a typical NEA can be reset by YORP is poorly understood.  Kaasalainen et al. (2007) determined that 1862 Apollo ($a=1.47$ AU, $D=1.4$ km) is spun up by YORP on a characteristic timescale $t_{\\rm YORP} = \\omega/(d\\omega/dt) \\sim 2.6$ My, where $\\omega$ is the spin rate. Starting from its current 3-hour spin period, 1862 Apollo is thus expected to be spun up to a $\\sim$2-hour period (and start shedding mass) in $\\sim$2 My. This timescale is probably at least slightly longer than the one on which the surface of 1862 Apollo should be reset by planetary encounters indicating that the YORP effect can be ignored for 1862 Apollo.  Since $t_{\\rm YORP} \\propto D^2 a^2 \\sqrt{1-e^2}$ (e.g., Nesvorn\\'y et al. 2007), however, the YORP effect can become more important than planetary encounters for small NEAs that orbit closer to the Sun than 1862 Apollo. For example, a sub-km NEA with $a<1$ AU can have $t_{\\rm YORP}$ several times shorter than 1862 Apollo. We therefore speculate that the YORP effect can contribute to the observed excess of Q-type NEAs in these low-$a$ orbits. A detailed analysis of this problem goes beyond the scope of this paper."
    },
    "1005/1005.1653_arXiv.txt": {
        "abstract": "We examine the clustering properties of a population of quasars drawn from fully hydrodynamic cosmological simulations that directly follow black hole growth.  We find that the black hole correlation function is best described by two distinct components: contributions from BH pairs occupying the same dark matter halo ('1-halo term', $\\xi_{\\rm{BH,1h}}$) which dominate at scales below $\\sim 300 \\: \\rm{kpc \\: h^{-1}}$, and contributions from BHs occupying separate halos ('2-halo term', $\\xi_{\\rm{BH,2h}}$ ) which dominate at larger scales.  From the 2-halo BH term we find a typical host halo mass for faint-end quasars (those probed in our simulation volumes) ranging from $M\\sim 10^{11}$ to a few $10^{12} M_{\\odot}$ from $z=5$ to $z=1$ respectively (consistent with the mean halo host mass). The BH correlation function shows a luminosity dependence as a function of redshift, though weak enough to be consistent with observational constraints.  At small scales, the high resolution of our simulations allows us to probe the 1-halo clustering in detail, finding that $\\xi_{\\rm{BH,1h}}$ follows an approximate power law, lacking the characteristic decrease in slope at small scales found in 1-halo terms for galaxies and dark matter.  We show that this difference is a direct result of a boost in the small-scale quasar bias caused by galaxies hosting multiple quasars (1-subhalo term) following a merger event, typically between a large central subgroup and a smaller, satellite subgroup hosting a relatively small black hole. We show that our predicted small-scale excess caused by such mergers is in good agreement with both the slope and amplitude indicated by recent small-scale measurements.  Finally, we note the excess to be a strong function of halo mass, such that the observed excess is well matched by the multiple black holes of intermediate mass ($10^7-10^8$ M$_{\\odot}$) found in hosts of $M \\sim 4-8 \\times 10^{11} M_\\odot$, a range well probed by our simulations. ",
        "introduction": "With supermassive black holes being found at the centre of most galaxies \\citep{1995ARA&A..33..581K}, interest in quasars has increased significantly, with substantial investigation into fundamental relations between black hole masses and their host galaxies' properties \\citep{1998AJ....115.2285M, 2000ApJ...539L...9F, 2000ApJ...539L..13G, 2002ApJ...574..740T, 2007ApJ...655...77G}.  In addition to these relations, statistical studies of the spatial clustering of quasars provide the potential to better understand the relation between quasars, their hosts and the underlying dark matter distribution, as well as estimate quasar lifetimes \\citep[see, e.g.,][]{HaimanHui2001, MartiniWeinberg2001} across a relatively large range of redshift. For example, strong clustering would suggest quasars should reside in massive groups. If so, they should be rare and in order to reproduce the quasar luminosity density, they must have long lifetimes. Conversely, low correlation would suggest more common quasars, and thus shorter quasar lifetimes.  Early studies of quasar clustering produced varying results for the clustering amplitude, with no clear agreement on overall evolution with redshift, some suggesting minimal or decreasing clustering evolution \\citep{MoFang1993, CroomShanks1996}, while others found an increase in clustering with redshift \\citep{Kundic1997, LaFranca1998}. These findings were generally poorly constrained due to the small sizes of available quasar samples. With the emergence of large scale surveys such as Sloan Digital Sky Survey \\citep{2000AJ....120.1579Y} and the Two-degree Field QSO Redshift Survey \\citep{2002MNRAS.333..279L}, substantially larger catalogs have been compiled, permitting more detailed investigation into the clustering properties of quasars, and many recent studies have been made into this area \\citep[e.g.][]{LaFranca1998, Porciani2004, Croom2005, Shen2007, Myers2007, daAngela2008, Shen2009, Ross2009}. These recent studies have found evidence for an increase in clustering amplitude with redshift \\citep{LaFranca1998, Porciani2004}, primarily for $z>2$, in agreement with predictions from simulations \\citep[see, e.g.][]{Bonoli2009, Croton2009}.  In addition to overall evolution, the luminosity dependence (if any) of large-scale clustering can provide significant insight into what quasar populations dominate different luminosity ranges. For example, the model of \\citet{2005ApJ...630..705H, 2005ApJ...630..716H, 2005ApJ...632...81H, 2005ApJ...625L..71H, 2006ApJS..163....1H} suggests that bright and faint quasars are similar objects which are observed at different phases of their lifetimes, rather than being fundamentally different populations of quasars (as simpler, 'on-off' models assume). This model would suggest that both bright and faint quasars should populate similar halos. Thus, while there may be some correlation between peak luminosity and host halo mass, clustering dependence on instantaneous luminosity should be relatively weak, particularly when compared to more traditional 'on-off' models of quasar luminosity \\citep{Lidz2006}. Recent observational studies have generally found a lack of luminosity dependence in the correlation function \\citep[see, e.g.,][]{Croom2005, Myers2007, daAngela2008}, though \\citet{Shen2009} found evidence for some, though weak, luminosity dependence.  Several semi-analytic models have also been used, finding differing luminosity dependences, such as a significant dependence for sufficient luminosity ranges, but limited when considering only luminosities probed by observation \\citep{Bonoli2009}, or weak dependence at low redshift ($z<1$), but stronger at higher redshift \\citep{Croton2009}.  In addition to large scale behavior, the possibility of excess quasar clustering on very small scales has arisen in several recent studies.  While some observed quasar pairs are believed to be the result of gravitationally lensed quasars, it has been proposed that others may be physically distinct quasar binaries, which would suggest quasars cluster much more strongly on small scales than extrapolation of large scale clustering would imply \\citep{Djorgovski1991, Hewett1998, Kochanek1999, Mortlock1999}, suggesting a connection between galaxy mergers and quasar activity \\citep[see, e.g.][]{Kochanek1999}.  However, investigating the smallest scale clustering has typically been problematic due to observational limitations (such as fiber collisions preventing small-separation pairs from being distinguished as distinct objects) and sample sizes insufficient for probing the smallest scales, where quasar pairs are rare.  There have been several studies probing clustering at sub-Mpc scales, generally finding no excess clustering relative to an extrapolation of the large-scale clustering behavior \\citep[see, e.g.][]{Shen2009b, Padmanabhan2009}.  However, these studies have been limited to scales above 100 $\\rm{kpc \\: h^{-1}}$, while several recent studies have managed to probe even smaller scales, where they do indeed find a significant excess \\citep{Hennawi2006, Myers2007II, Myers2008}.   In particular, \\citet{Hennawi2006} studied binary quasars from SDSS and 2dF Quasar Survey to compute quasar clustering for scales as small as 20 $\\rm{kpc} \\: h^{-1}$ (comoving), and found significant excess clustering relative to the large scale extrapolation (by an order of magnitude at comoving scales below 100 $\\rm{kpc} \\: h^{-1}$, and growing stronger with decreasing scale).  This excess implies that the quasars are more strongly clustered than galaxies at these small scales, supporting the theory that quasar activity is triggered by galaxy interactions.  Using the quasar sample from \\citet{Myers2007}, \\citet{Myers2007II} found only a slight excess in small-scale clustering, and put an upper limit for the excess at a factor of 4.3\\underline{+}1.3 for physical scales of $\\sim 28 \\: \\rm{kpc} \\: h^{-1}$.  They suggest that the significantly larger excess of \\citet{Hennawi2006} is a result of a selection effect, possibly due to studies tending to target tracers of the Ly$\\alpha$ forest, causing a bias toward $z > 2$, which may be more highly clustered. \\citet{Myers2008} used a complete spectroscopic sample of quasars over physical scales of 23.7-29.9 $\\rm{kpc} \\: h^{-1}$ from SDSS to find an excess clustering factor of $\\sim$4, consistent with the upper limit of \\citet{Myers2007II}, which, while 2$\\sigma$ below the excess found by \\citet{Hennawi2006}, nonetheless supports the general finding of a clustering excess which may be a result of galaxy interactions.  In this paper, we use cosmological hydrodynamic simulations which directly model the growth, accretion, and feedback processes of black holes to investigate the properties and underlying causes of black hole clustering. Although the simulation volume limits our analysis to black hole luminosities and host group masses below those typically studied, the self-consistent modeling of black holes allows us to study the clustering behavior without post-processing models.  Additionally, the high resolution allows us to investigate clustering behavior at extremely small scales, well below those studied with semi-analytic models, thereby providing a means of using simulations to investigate the observed small-scale excess for the first time, and provide a physical explanation for the underlying cause.  In Section 2 we describe the numerical modeling for the black holes formation and accretion (Section 2.1) the simulation parameters used (Section 2.2), the details of the subgroup finder (Section 2.3) and our method of calculating correlation functions (Section 2.4).  In Section 3 we investigate the quasar clustering properties at both large and small scales, and we summarize our results in Section 4.  \\begin{figure*} \\centering \\includegraphics[width=14cm]{plots/temp4.eps} \\caption{An example of the distribution of black holes in the simulations: The same slice ($2 \\: \\rm{Mpc} \\: h^{-1}$ thick) through the D6 simulation at z=1,2,3,4. The positions of black holes in different luminosities bins ($L < 10^8 L_\\odot$ - Orange; $10^8 L_\\odot < L < 10^9 L_\\odot$ - Pink; $10^9 L_\\odot < L < 10^{10} L_\\odot$ - Blue; $L > 10^{10} L_\\odot$ - Green.) are plotted on top of the gas density distribution (shown in the the gray scale).} \\label{simulationslice} \\end{figure*} ",
        "conclusions": "In this paper we have investigated the clustering of black holes within hydrodynamic cosmological simulations, its redshift evolution, luminosity dependence, and particularly the small-scale behavior.  We have shown that the large scale clustering of black holes traces that of the galaxies within their host groups, and provides a predictor of the typical host mass, which for our simulations is found to be on the order of a few $10^{11} M_\\odot$.  Although well below the typically found masses of $\\sim 2 \\times 10^{12}-10^{13} M_\\odot$ \\citep{Lidz2006,Ross2009,Bonoli2009, Shen2009}, this is consistent with our limited simulation volumes which can only follow the growth of the faint-end of quasar population (DeGraf et al. 2010), and cannot follow formation of such massive groups. The typical host group mass shows some evolution with redshift, most significant below $z \\sim 3$, where typical host masses increase by up to a factor 10 (at $z=1$). This low-redshift increase is distinctly luminosity dependent, with the more luminous sources ($L_{\\rm{BH}} > 10^9 L_\\odot$) undergoing the most substantial increase in typical host mass.  Overall the evolution of clustering with redshift and luminosity is minor and consistent with current observational constraints (albeit in low luminosity populations this is yet to be fully constrained). The relatively weak dependence found in our simulations is consistent with the complex lightcurves we derive from our direct modeling in which quasar luminosities vary over relatively short timescales for a given host (as regulated by hydrodynamical processes).  This is also consistent with the models of \\citet{Lidz2006}.  In addition to the large-scale clustering (the 2-halo regime), our simulations allow us to study the small scale clustering (the 1-halo term) of $\\xi_{\\rm{BH}}$.  We found that $\\xi_{\\rm{BH,1h}}$ follows a power law behavior all the way to the smallest scales. The clustering of black holes at small scale is unlike that of galaxies (or dark matter).  We showed that the 1-halo BH term can be subdivided into two components: 1-subhalo and 2-subhalo. The 1-subhalo term, $\\xi_{\\rm{subgroup,1h}}$, represents the clustering of BHs within a galaxy and 2-subhalo that of BHs occupying different galaxies.  We have shown that the 1-subhalo is the one that provides the power law behavior, indicating that galaxies do contain multiple black holes as a result of mergers. These galaxies tend to be the central galaxy within relatively large groups (for our simulation), generally hosting at most a single massive BH with one or more smaller BHs, likely as a result of smaller satellite galaxies merging with the large, central galaxy within the group.  In the absence of these multiply-occupied galaxies, $\\xi_{\\rm{BH,1h}}$ and $\\xi_{\\rm{subgroup,1h}}$ exhibit very close agreement, but the inclusion of these merger remnants causes a significant boost in the small-scale BH clustering.  This merger-based boost is most significant at low redshift, where typical group size is largest, though we find it in sufficiently massive groups at all redshifts.  Though observational limitations make observing these scales difficult, several recent studies have found a small scale excess at scales below $\\sim 100 \\: \\rm{kpc} \\: h^{-1}$ \\citep{Hennawi2006, Myers2008}. The observed excess is in remarkable agreement to the one predicted by our simulations coming from groups approaching $10^{12} M_\\odot$, which host mostly intermediate size black holes. This suggests that multiple black holes co-occupying a subgroup at low redshifts are likely faint(ish) AGNs hosted in Milky Way size halos that have recently undergone merging.  We also note that galaxies hosting multiple AGN \\citep{Komossa2003, Gerke2007, Barth2008, Comerford2009b} or inspiralling supermassive black holes \\citep{Comerford2009} have been found in recent studies, further supporting our conclusion of multiply-occupied subgroups. Although we leave more detailed investigation of the small scale BH pairs in our simulations (particularly with regard to the luminosities of inspiralling black holes) for a future work, we note that our finding that multiply-occupied galaxies tend to host a single massive BH with one or more small BHs appears to be in keeping with the observation that most of the inspiralling BH pairs power only a single AGN \\citep{Comerford2009}.  Given that, our agreement in small-scale merger-induced boost certainly reinforces the importance of galaxy mergers on the evolution of supermassive black holes. We also note this small-scale excess' sensitivity to the host mass suggests that future small-scale studies may provide a means to constrain the typical mass of merger events between galaxies hosting black holes, with current observational data combined with our simulations suggesting groups with typical masses comparable to those probed in our simulations (from a few $10^{11} M_\\odot$ to $10^{12} M_\\odot$) produce the multiply-occupied galaxies underlying the observed small scale excess.  We would like to point out however that there are several aspect of our modeling approach, including numerical issues, in the simulations that potentially affect our results on the small-scale clustering.  We have a very simplistic prescription to determine how BHs merge with one another (imposed by the limits on the resolution that can be achieved in these cosmological boxes). The current prescription has a BH pair merge when BHs are separated by less than their smoothing length and if the BHs relative velocity is small compared to the local sound speed. Changes to this prescription could accelerate (postpone) BH mergers, which would result in a suppression (increase) of our small scale clustering signals. It would be desirable to compare our results with other simulations which implement different prescriptions, or in the future to include more direct physical modeling of this region in higher resolution simulations.  However, neither of these are currently possible.  A numerical issue that may affect the results of our one-halo term is that black holes need to be fixed to potential minima (calculated among the neighboring particles within the smoothing length used for the accretion model) in order to avoid them leaving their subhalo due to numerical N-body noise (and the fact that dynamical friction is hard to calculate for sink particles).  However, in some instances this may cause a BH particle in a small subhalo in orbit in a bigger group to 'hop' to the potential minimum of the larger group. This effect may be exacerbated in situations where the small subhalo may be stripped of gas by infalling into a larger one. These effects could artificially increase the number of BHs within large, central halos, thereby boosting small scale clustering. However, when we measure what fraction of BHs appear to 'hop' into the center of groups experiencing an unexpected jump in their position, we find that it is only $\\sim 1-2\\%$.  Future simulations and comparison amongst different approaches (once they become available) should of course attempt to characterize these effects more specifically.  We further note however, that observational studies have indeed found cases of galaxies hosting multiple BHs \\citep{Comerford2009}, so the existence of a one-subhalo term is expected. Additionally, as seen in Figure \\ref{projectedxi}, the projected clustering of subgroups has a fundamentally different form than the observed quasar clustering.  Thus the BHs cannot simply trace their host subgroups/galaxies and still produce the observed small scale excess, but rather a significant one-subhalo term is required to produce the small scale power law behavior.  In future work we also plan to simulate larger volumes (which we are starting to be feasible with the most advanced technology) to allow us to study clustering of AGN at larger (mass and length) scales while simultaneously investigating luminosity dependence for brighter sources more directly comparable to current and upcoming observational data, as well as providing increased statistics for the small scale clustering."
    },
    "1005/1005.3656_arXiv.txt": {
        "abstract": "We present near-UV transmission spectroscopy of the highly irradiated transiting exoplanet WASP-12b, obtained with the Cosmic Origins Spectrograph (COS) on the Hubble Space Telescope (HST). The spectra cover three distinct wavelength ranges: NUVA (2539--2580\\,\\AA); NUVB (2655--2696\\,\\AA); and NUVC (2770--2811\\,\\AA). Three independent methods all reveal enhanced transit depths attributable to absorption by resonance lines of metals in the exosphere of WASP-12b. Light curves of total counts in the NUVA and NUVC wavelength ranges show a detection at a 2.5$\\sigma$ level. We detect extra absorption in the \\ion{Mg}{2} $\\lambda\\lambda$2800 resonance line cores at the 2.8$\\sigma$ level. The NUVA, NUVB and NUVC light curves imply effective radii of 2.69$\\pm$0.24\\,R$_J$, 2.18$\\pm$0.18\\,R$_J$, and 2.66$\\pm$0.22\\,R$_J$ respectively, suggesting the planet is surrounded by an absorbing cloud which overfills the Roche lobe. We detect enhanced transit depths at the wavelengths of resonance lines of neutral sodium, tin and manganese, and at singly ionised ytterbium, scandium, manganese, aluminum, vanadium and magnesium. We also find the statistically expected number of anomalous transit depths at wavelengths not associated with any known resonance line. Our data are limited by photon noise, but taken as a whole the results are strong evidence for an extended absorbing exosphere surrounding the planet. The NUVA data exhibits an early ingress, contrary to model expectations; we speculate this could be due to the presence of a disk of previously stripped material. ",
        "introduction": "Observations of the transiting extrasolar planets HD209458b and HD189733b revealed an enhanced transit depth at the wavelengths of several UV resonance lines \\citep{vidal03,vidal04,lecavelier10}. These UV lines from the ground state are sensitive probes of the presence of atomic and ionic species. Their presence enhanced the effective radius of the planet during transit, implying the planet is surrounded by an extended cloud of size comparable to or larger than its Roche lobe \\citep{vidal03,vidal04,jaffel07,vidal08}. This was attributed to a hydrodynamic `blow-off' of the planet's outer atmosphere caused by the intense irradiation suffered by this hot Jupiter exoplanet. An alternative explanation in which the planet is surrounded by a cloud of energetic neutral atoms caused by interactions with the host star's stellar wind has, however, been suggested \\citep{ena08,ena10}. WASP-12b is one of the hottest and most irradiated transiting exoplanets and orbits extremely close to a late F-type host star \\citep{hebb2009}. WASP-12b is, therefore, an attractive target to explore the properties of the phenomenon observed in HD209458b, and might yield evidence distinguishing between the suggested underlying causes.  The initial UV observations of HD209458b were in the far UV around the Ly\\,$\\alpha$ emission line. The abundance of hydrogen makes this an attractive line to observe, but the temporal and spatial variability of stellar Ly\\,$\\alpha$ emission is a highly undesirable complicating factor. For this reason, and to obtain better signal to noise, we observed WASP-12 in the near-UV where there are many other resonance lines \\citep{morton1991,morton2000}, including the very strong \\ion{Mg}{2} UV resonance lines. This work became possible with the installation of the Cosmic Origins Spectrograph (COS) on the Hubble Space Telescope (HST) reinstating and enhancing our capabilities for UV spectroscopy. ",
        "conclusions": "\\label{sec:disc} We have performed three independent analyses, each of which suggests absorption in the resonance lines of metals from an extended atmosphere surrounding the transiting planet WASP-12b. In Section~\\ref{charb} we found a deeper transit  in the core region of the \\ion{Mg}{2} doublet at the $2.8\\sigma$ level.  In Sect.~\\ref{sec_lightcurve}, the transit depths in the NUVA, NUVB, and NUVC wavelength ranges respectively imply effective planet radii of 2.69$\\pm$0.24\\,R$_J$, 2.18$\\pm$0.18\\,R$_J$, and 2.66$\\pm$0.22\\,R$_J$. WASP-12b's optical radius is $R_{\\rm P} = 1.79 \\pm 0.09 \\,R_J$ while the mean Roche lobe radius is 2.36\\,R$_J$ using Paczy{\\'n}ski's (1971) prescription.  Table~\\ref{wavelengths} shows that we detect enhanced transit depths at the wavelengths of resonance lines of neutral sodium, tin and manganese, and at singly ionised ytterbium, scandium, manganese, aluminum, vanadium and magnesium. Finally we detect an enhanced transit depth within $0.12$\\AA\\ of a resonance line of doubly ionised europium. We also find the statistically expected number of anomalous transit depths at wavelengths not associated with any known resonance line.  Taken as a whole, these results constitute compelling evidence that WASP-12b is surrounded by an exosphere which over-fills the planet's Roche lobe, confirming predictions by \\citet{li2010}. This exosphere is likely composed of a number of elements/ions, including probably \\ion{Na}{1}, \\ion{Mg}{1}, \\ion{Mg}{2}, \\ion{Al}{1}, \\ion{Sc}{2}, \\ion{Mn}{2}, \\ion{Fe}{1}, and \\ion{Co}{1}. The phenomenon found in HD209458b \\citep{vidal03,vidal08} probably occurs generally for hot Jupiter exoplanets. By analogy with HD209458b, and as WASP-12b and its host star are almost certainly predominantly composed of hydrogen, we expect that this exosphere is hydrogen rich.  Models by \\citet{yelle} suggest that elements other then H and He should not be present in the upper atmosphere due to the low vertical mixing rate, but this takes Jupiter as the starting point. WASP-12b is extremely close to the host star and consequently the stellar irradiation and tidal effects could induce prodigious mixing, affecting the chemistry of the planet atmosphere. Our detections of several metallic elements and/or ions is certainly consistent with a metal-rich atmosphere for WASP-12b.  The most surprising result is provided by the juxtaposition of our data with the optical ephemeris. We took contemporaneous optical photometry with OU-OAM PIRATE \\citep{klb2009} which showed the ephemeris of \\citet{hebb2009} remains accurate. Figure~\\ref{lightcurve} shows the NUVA transit has an early ingress and an egress consistent with the optical ephemeris. In contrast, naive momentum considerations and hydrodynamic simulations would instead suggest that the effect of a diffuse cloud surrounding the planet would be to smear and delay egress while ingress is relatively unaffected, see e.g. Fig.~1 and 2 of \\citet{schneiter2007}.  In detail, the shape of the diffuse cloud may well be element/ion dependent since different elements/ions behave differently in the presence of strong radiation pressure. This can explain why we observe different transit shapes in the NUVA region and the other regions. As Fig.~\\ref{spectra} shows, the stellar spectrum in the NUVA region is strongly absorbed by a plethora of lines, dominated those of neutral elements.  The NUVC region is also strongly absorbed in the stellar photosphere but predominantly from the \\ion{Mg}{2} doublet. It is presumably the cumulative absorption from many relatively weak spectral lines in the planet's exosphere which creates the excess transit depth in the NUVA region, while Table~\\ref{wavelengths} and Eq.~\\ref{mg2depth} demonstrate that planet's absorption in the NUVC region is associated with the \\ion{Mg}{2} doublet. The \\ion{Mg}{2} ion will experience different forces to neutral atoms in an environment where there is certainly a strong radiation field, and  strong and varying large-scale magnetic fields are also likely. The NUVB light curve is least deviant from the optical transit, and this is consistent with the relative dearth of strongly absorbing lines in this spectral window, c.f. Fig.~\\ref{spectra}.  We do not have any detailed explanation for the observed early ingress in NUVA, but we speculate the effect could be produced  if material is lost from the planet exosphere and forms a diffuse ring or torus around the star enveloping the planet's orbital path, as models suggest \\citep{li2010}. The orbital motion of the planet through this medium might compress the material in front of it. This could increase the opacity of the medium through which the star is viewed immediately before first contact. A void in the medium might be expected to form behind the planet, and consequently the egress is relatively unaffected by the diffuse ring.  Our observations demonstrate that COS spectroscopy of transiting exoplanets has the potential to detect many species via transmission spectroscopy, and to measure velocities and deduce spatial distributions. There are now about 40 known transiting exoplanets with orbital periods shorter than that of HD209458b. Many of these transit stars significantly brighter than WASP-12b. COS spectroscopy of brighter examples will allow us probe the exosphere species-by-species examining their density, velocity and spatial distributions. This detailed information should allow us to determine whether these planets really are being photo-evaporated by their host stars, and, if so, to empirically deduce the mass loss rate. We encourage detailed element/ion dependent modeling of the exosphere in the highly irradiated environment of  WASP-12b and similar systems, and observations of other similar extrasolar planets. There is a rich new parameter space to explore!"
    },
    "1005/1005.3642_arXiv.txt": {
        "abstract": "{The Ultra Luminous InfraRed Galaxy (ULIRG) Mrk\\,231 reveals up to seven rotational lines of water (H$_2$O) in emission, including a very high-lying ($E_{\\mathrm{upper}}=640$ K) line detected at a 4$\\sigma$ level, within the Herschel/SPIRE wavelength range ($190 < \\lambda (\\mathrm{\\mu m})< 640$), whereas PACS observations show one H$_2$O line at 78 $\\mu$m in absorption, as found for other H$_2$O lines previously detected by ISO. The absorption/emission dichotomy is caused by the pumping of the rotational levels by far-infrared radiation emitted by dust, and subsequent relaxation through lines at longer wavelengths, which allows us to estimate both the column density of H$_2$O and the general characteristics of the underlying far-infrared continuum source. Radiative transfer models including excitation through both absorption of far-infrared radiation emitted by dust and collisions are used to calculate the equilibrium level populations of H$_2$O and the corresponding line fluxes. The highest-lying H$_2$O lines detected in emission, with levels at $300-640$ K above the ground state, indicate that the source of far-infrared radiation responsible for the pumping is compact (radius$=110-180$ pc) and warm ($T_{\\mathrm{dust}}=85-95$ K), accounting for at least $45$\\% of the bolometric luminosity. The high column density, $N(\\mathrm{H_2O})\\sim5\\times10^{17}$ cm$^{-2}$, found in this nuclear component, is most probably the consequence of shocks/cosmic rays, an XDR chemistry, and/or an ``undepleted chemistry'' where grain mantles are evaporated. A more extended region, presumably the inner region of the 1-kpc disk observed in other molecular species, could contribute to the flux observed in low-lying H$_2$O lines through dense hot cores, and/or shocks. The H$_2$O 78 $\\mu$m line observed with PACS shows hints of a blue-shifted wing seen in absorption, possibly indicating the occurrence of H$_2$O in the prominent outflow detected  in OH (Fischer et al., this volume). Additional PACS/HIFI observations of H$_2$O lines are required to constrain the kinematics of the nuclear component, as well as the distribution of H$_2$O relative to the warm dust.} ",
        "introduction": "One key question in the study of composite infrared (IR) merging galaxies and quasi-stellar objects (QSOs) is what fraction of their luminosity is generated in the nuclear region ($<200$ pc) associated with the Active Galactic Nucleus (AGN) and a possible extreme nuclear starburst, and what fraction arises from a more extended kpc-scale starburst \\citep[e.g.][]{arm07,vei09}. The ULIRG Markarian 231 (Mrk 231) is the most luminous ($L\\sim4\\times10^{12}$ \\Lsun) galaxy in the local Universe ($z<0.1$), and thus provides a unique template for such studies. Since the bulk of the luminosity in ULIRGs arises at far-IR wavelengths, where sub-arc-second resolution observations are not available, an alternative technique is required to constrain the compactness of the far-IR emission and its physical origin.  In a previous work based on observations with the \\emph{Infrared Space Observatory (ISO)}, \\citet[][hereafter G-A08]{gon08} have argued  that the observation of molecular species such as OH and H$_2$O at far-IR wavelengths is ideal for such a purpose, because their high-lying rotational levels are pumped through absorption of far-IR radiation and the observable excitation is then sensitive to the far-IR radiation density that in turn depends on the compactness of the far-IR continuum source. In addition, these molecular observations shed light on the dominant chemistry in those nuclear regions. G-A08 reported the ISO detection of 3 high-lying H$_2$O lines, relevant upper limits over the entire ISO spectrum, and also high-lying OH lines, indicating the ocurrence of a compact-luminous far-IR component.   With their high sensitivity, spectral resolution, and wavelength coverage, the \\emph{Herschel} \\citep{pil10} instruments are ideal for extending our previous study to additional key lines in the far-IR/submillimeter. As part of the HerCULES Key Programme \\citep[see][this volume, hereafter vdW10]{vdw10}, we report in this Letter the Herschel SPIRE/PACS \\citep{gri10,pog10} detection and first analysis of several H$_2$O lines in Mrk 231, which supports the conclusions of G-A08 and gives additional clues to the origin of H$_2$O in this ULIRG. We adopt a distance to Mrk 231 of 192 Mpc ($H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.73$, and $z=0.04217$). ",
        "conclusions": "The extreme nature of the nuclear region in Mrk 231 is well illustrated by comparing its SPIRE spectrum with that of the Orion Bar \\citep[][this issue]{hab10}, the prototypical Galactic PDR. The Orion Bar spectrum shows CO lines a factor of $\\gtrsim50$ stronger than the H$_2$O lines, while in Mrk 231 the H$_2$O and CO lines have comparable strengths (vdW10). This contrast will be still higher in the nuclear region, provided that a significant fraction of the CO emission in Mrk 231 arises from a more extended region. Thus the H$_2$O-to-CO line intensity ratios in the SPIRE wavelength range are an excellent diagnostic of extragalactic compact/warm far-IR continuum sources with unusually high amounts of H$_2$O.  The above comparison also indicates that the nuclear region of Mrk 231 cannot be interpreted as an ensemble of classical PDRs. Three main scenarios are proposed to explain such high amounts of H$_2$O: $(i)$ widespread shocks/cosmic rays: although the H$_2$O lines peak around the systemic velocity, outflows of $\\sim100$ km s$^{-1}$ are not ruled out by our data, and indeed some indications in the H$_2$O line shapes of systematic motions have been found; an enhanced cosmic ray flux could also have an important impact on the nuclear chemistry. $(ii)$ XDR chemistry: our derived H$_2$O abundance of $\\sim10^{-6}$ is in very good agreement with XDR model results by \\citet[][their Fig. 3, Model 3]{mei05}, as well as with our preliminary estimate of the H$_2$O spatial distribution. $(iii)$ An undepleted chemistry, where H$_2$O that formed on grain mantles is released into the gas phase, as in Galactic hot cores; in support of this scenario, the derived $T_{\\mathrm{dust}}$ in $W_C$ is close to the evaporation temperature of solid H$_2$O. All three scenarios are probably taking place, and the identification of the dominant process requires a multi-species analysis.  The nuclear region traced by the high-lying H$_2$O lines has a size similar to the nuclear disk (or outflow) observed at radio wavelengths and H I 21 cm by \\citet[][their Figs. 3 \\& 7]{car98}, suggesting a close physical correspondence. From $H_C$ and $W_C$, the nuclear surface brightness ($\\sim1.5\\times10^{13}$ \\Lsun\\ kpc$^{-2}$) exceeds the highest values attained in starburst on spatial scales $\\gtrsim100$ pc \\citep{meu97,dav07}, while the nuclear luminosity-to-mass ratio ($L/M\\sim4\\times10^3$ \\Lsun/\\Msun) exceeds the limit for a starburst estimated by \\citet{sco03}. From near-IR data, \\citet{dav07} estimated a starburst luminosity from a similarly sized region of $\\lesssim7\\times10^{11}$ \\Lsun; according to the joint luminosity of our $H_C$ and $W_C$, the AGN would account for at least 50\\% of the output power in Mrk 231."
    },
    "1005/1005.1647_arXiv.txt": {
        "abstract": "\\noindent Numerical simulations show that box-shaped bulges of edge-on galaxies are not bulges: they are bars \\hbox{seen side-on.}  Therefore the two components that are seen in edge-on Sb galaxies such as NGC 4565 are a disk and a bar.  But face-on SBb galaxies always show a disk, a bar, and a (pseudo)bulge.~Where is the (pseudo)bulge in NGC{\\thinspace}4565?  We use archival {\\it Hubble Space Telescope\\/} $H$-band images and {\\it Spitzer Space Telescope\\/} 3.6 $\\mu$m wavelength images, both calibrated to 2MASS $K_s$ band, to penetrate the prominent dust lane in NGC 4565.  We find a high surface brightness, central stellar component that is clearly distinct from the boxy bar and from the disk.  Its brightness profile is a S\\'{e}rsic function with index $n = 1.55 \\pm 0.07$ along the major axis and 1.33$\\pm$0.12 along the minor axis.  Therefore it is a pseudobulge.  It  is much less luminous than the boxy bar, so the true pseudobulge-to-total luminosity ratio of the galaxy is  $PB/T = 0.06 \\pm 0.01$, much less than the previously believed value of $B/T = 0.4$ for the ``boxy bulge''.  We infer that published $B/T$ luminosity ratios of edge-on galaxies with boxy bulges have been overestimated.  Therefore, more galaxies than we thought contain little or no evidence of a merger-built classical bulge.  From a formation point of view, NGC 4565 is a giant, pure-disk galaxy.  This presents a challenge to our picture of galaxy formation by hierarchical clustering: it is difficult to grow galaxies as big as NGC 4565 without also making big classical bulges. ",
        "introduction": "\\pretolerance=15000  \\tolerance=15000   Figure 1 compares the prototypical Sb galaxies NGC 3351 and NGC 4565.  NGC 3351 is more nearly face-on and shows three main components -- a bulge, a bar,~and~a~disk.  In contrast, NGC 4565 is edge-on; it shows only two components, a box-shaped bulge and disk.  As long as we thought that boxy structure was a secondary property of normal bulges, a galaxy morphologist (Sandage 1961) would just use the bulge-to-total luminosity ratio $B/T \\simeq 0.4$ (Simien \\& de Vaucouleurs 1986) to classify NGC 4565 as an Sb.  However, we now know that \\phantom{00000000}  \\vskip 230pt   \\centerline{\\null} \\vskip -6pt  \\noindent   ``boxy bulges'' are not bulges at all; rather, they are edge-on bars (Combes \\& Sanders 1981).  So the SBb galaxy NGC 3351 shows a disk, a bar, and a bulge, but the SBb galaxy NGC 4565 shows only a disk and a bar.  Where is the bulge in NGC 4565?  Boxy bulges like that in NGC 4565 are a fundamental feature of edge-on disk galaxies (Sandage 1961; Buta \\etal 2007).  Their identification as edge-on bars is a well known result. The Combes \\& Sanders (1981) $N$-body model demonstration that bars heat themselves vertically by a combination of buckling instabilities and resonant star scattering has been confirmed \\phantom{00000000}  \\clearpage  \\noindent and extended many times (Combes et al.~1990; Pfenniger~\\& Norman 1990; Pfenniger \\& Friedli 1991; Raha et al.~1991; Athanassoula \\& Misiriotis 2002; Athanassoula 2005, Shen et al.\\ 2010, and many others). Cylindrical rotation is observed in $N$-body bars and in boxy bulges (Kormendy \\& Illingworth 1982; Jarvis 1990; Shaw, Wilkinson, \\& Carter 1993; Bettoni \\& Galletta 1994; Fisher \\etal 1994; D'Onofrio et al.~1999; Falc\\'on-Barroso et al.~2004; Howard \\etal 2008) but not in classical, elliptical-galaxy-like bulges (Illingworth \\&  Schechter 1982; Kormendy \\& Illingworth 1982; Binney \\etal 1990; Verolme et al.~2002; Copin \\etal 2004; Emsellem \\etal 2004); this further cements the connection between boxy bulges and edge-on bars. Finally, a splitting of gas rotation velocities in edge-on boxy bulges (a ``figure 8'' shape of spectral emission lines) also is a robust signature of gas flow in an edge-on bar (Kuijken \\& Merrifield 1995; Merrifield 1996; Merrifield \\& Kuijken 1999; Bureau \\& Freeman 1999). There is little doubt that the ``boxy bulge'' of NGC 4565 is not the real bulge of the galaxy.  Does the galaxy contain a bulge at all?  That is, does it contain a dense, central component that we would identify as a bulge in addition to the bar if the galaxy were seen face-on?  We care for two reasons, both involving galaxy formation:  {\\it Background:} Early galaxy evolution was dominated by hierarchical gravitational clustering of density fluctuations that resulted in galaxy collisions and mergers \\hbox{(White \\& Rees 1978);} these scrambled disks into ellipticals (Toomre~1977).~Enormous energy has been invested in studying hierarchical clustering; there is little danger that the picture is fundamentally wrong (Binney 2004).  However, it is incomplete. Recent work has established that hierarchical clustering is gradually giving way to a complementary suite of evolution processes that shape {\\it isolated\\/} galaxies.  They evolve by rearranging energy and angular momentum; one consequence is the growth of central components that masquerade as classical bulges but that, in general, formed slowly (``secularly'') out of disks (see Kormendy 1993; Kormendy \\& Kennicutt 2004 for reviews).  We call them ``pseudobulges'' to distinguish them from merger remnants.  They come in at least two varieties.   As reviewed above, ``boxy bulges'' are believed to be edge-on bars. Our Galaxy contains one (Dwek \\etal 1995).   Another variety is grown out of disk gas that was transported inward by nonaxisymmetries such as bars; we call them ``disky pseudobulges'' here, because they are often highly flattened, but we emphasize that they are not always flat (Kormendy 1993; Kormendy \\& Kennicutt 2004; \\S\\thinspace2 here).  Then:  {\\it Reason 1:} Confidence in our conclusion that the boxy center of NGC 4565 is an edge-on bar would be increased if we also observed a (pseudo)bulge as we do in face-on galaxies~(Fig.~1). As long as face-on and edge-on galaxies appear to show physical differences, we can't be sure that we understand them.  {\\it Reason 2:} If the box in NGC 4565 is a bar, then it is part of the disk and $B/T$ is smaller than we thought.  If in addition the galaxy contains a pseudobulge and not a hidden classical bulge, then $B/T$ is even smaller -- possibly zero.  This is hard to understand in the context of hierarchical clustering, which essentially always makes substantial bulges in giant galaxies (see, e.{\\thinspace}g., Peebles \\& Nusser 2010; Kormendy \\etal 2010). ",
        "conclusions": "The interpretation of boxy bulges in edge-on Sb galaxies as bars is more believable if we also find (pseudo)bulges like those associated with bars in face-on Sb galaxies.  Our discovery of a pseudobulge in NGC 4565 that is distinct from the boxy bar increases confidence in our picture of secular evolution.  Furthermore, $B/T$ ratios in edge-on galaxies with boxy bulges are smaller than previously believed.   In NGC 4565, the detection only of a central pseudobulge means that we see no sign of a major merger remnant.   Moreover, NGC 4565 rotates at 255 km s$^{-1}$ interior to the outer warp (Rupen 1991).   NGC{\\thinspace}4565 has grown very massive while remaining a pure-disk galaxy.  We do not know how this can happen in a Universe dominated by the dynamical violence of hierarchical clustering.  \\centerline{\\null} \\vskip -30pt \\centerline{\\null}"
    },
    "1005/1005.4062_arXiv.txt": {
        "abstract": "Recent high resolution observations of the Galactic center black hole allow for direct comparison with accretion disk simulations. We compare two-temperature synchrotron emission models from three dimensional, general relativistic magnetohydrodynamic simulations to millimeter observations of Sgr A*. Fits to very long baseline interferometry and spectral index measurements disfavor the monochromatic face-on black hole shadow models from our previous work. Inclination angles $\\le 20^\\circ$ are ruled out to $3\\sigma$. We estimate the inclination and position angles of the black hole, as well as the electron temperature of the accretion flow and the accretion rate, to be $i={50^\\circ}^{+35{}^\\circ}_{-15{}^\\circ}$, $\\xi={-23^\\circ}^{+97{}^\\circ}_{-22{}^\\circ}$, $T_e=(5.4 \\pm 3.0) \\times 10^{10}$K and $\\dot{M}=5^{+15}_{-2}\\times10^{-9} M_\\odot \\mathrm{yr}^{-1}$ respectively, with 90\\% confidence. The black hole shadow is unobscured in all best fit models, and may be detected by observations on baselines between Chile and California, Arizona or Mexico at $1.3$mm or $.87$mm either through direct sampling of the visibility amplitude or using closure phase information. Millimeter flaring behavior consistent with the observations is present in all viable models, and is caused by magnetic turbulence in the inner radii of the accretion flow. The variability at optically thin frequencies is strongly correlated with that in the accretion rate. The simulations provide a universal picture of the $1.3$mm emission region as a small region near the midplane in the inner radii of the accretion flow, which is roughly isothermal and has $\\nu/\\nu_c \\sim 1-20$, where $\\nu_c$ is the critical frequency for thermal synchrotron emission. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.1701_arXiv.txt": {
        "abstract": "We searched for ${\\rm H_{2}O}$ 6(1,6)-5(2,3) maser emission at 22.235 GHz from several Saturnian satellites with the Nobeyama 45m radio telescope  in May 2009.  Observations were made for Titan, Hyperion, Enceladus and Atlas, for which Pogrebenko {\\it et al.} (2009) had reported detections of water masers at 22.235 GHz, and in addition for Iapetus and  other inner satellites. We detected  no emission of the water maser line for all the satellites observed, % although sensitivities of our observations were comparable or even better than those of  Pogrebenko {\\it et al.}. % We infer that the water maser emission from the Saturnian system is extremely weak, or  sporadic in nature.  Monitoring over a long period and obtaining statistical results must be made for the further understanding of the water maser emission in the Saturnian system. ",
        "introduction": "Maser emissions are widely found in celestial objects such as dense cores of molecular clouds and circumstellar envelopes of  late-type stars \\citep{rei81}.  Masers have been used as  probes of gas with the H$_2$ number density of typically $10^4$--$10^{10}$ cm$^{-3}$. For solar system objects, several maser and laser phenomena have been found; $e.g.$, ${\\rm CO_{2}}$ (Venus and Mars: Mumma 1992) and ${\\rm OH}$ (for many comets: {\\it e.g.} Crovisier {\\it et al.} 2002). Each phenomenon would be induced by  different physical processes. While a thermal 22.235 GHz water line was possibly detected for comet Hale-Bopp (Bird {\\it et al.} 1997), the first detection of ${\\rm H_{2}O}$ maser in the solar system was reported at the catastrophic impact of comet Shoemaker-Levy9 and Jupiter (Cosmovici {\\it et al.} 1996). This report suggests that such an incident can induce collisional pumping for water masers. Recently, Pogrebenko {\\it et al.} (2009) (abbreviated as POG hereafter) have reported the detections of ${\\rm H_{2}O}$ masers from the Saturnian satellites (Titan, Hyperion, Enceladus and Atlas) with the Medicina 32m and Mets${\\rm \\ddot{a}}$hovi 14m telescopes. This is interesting because, unlike a temporal phenomenon such as the break-up and disruption of a comet, we can perform long period monitoring of ${\\rm H_{2}O}$ emission using ground-based telescopes, space telescopes ({\\it e.g.} Herschel Space Telescope) and spacecrafts.  So far, we do not have much knowledge about the maser mechanism in the solar system, although a lot of water maser phenomena are observationally studied for extra-solar objects. The combination of ground, space and in-situ observations would contribute to understand the nature of water maser emission if the presence in the Saturnian system is verified.  Therefore, we must accumulate more data of the water maser lines in the Saturnian system. In this letter, we report our trial of detecting water maser satellites with the 45m radio telescope at Nobeyama Radio Observatory (NRO). We observed the major Saturnian satellites for which POG reported the detections, and in addition, we observed a few other inner satellites. ",
        "conclusions": "Figure 1 shows the obtained spectra (upper two  and lower-left panels) and the illustration of the satellite positions (lower-left panel) at the time of observations for Titan, Hyperion and Iapetus. Figure 2 also shows the spectra (upper panels) and illustrations of positions (lower panels)  for Enceladus and Atlas. The position of a satellite $\\theta$ is defined as an angle between the line of sight towards Saturn and the line of a satellite and Saturn ($i.e.$ $\\theta=0$ when the object transits on the Saturn seen from the Earth). Table 2 shows a typical 1$\\sigma$ upper limit of the signal considering the factor of a satellite position inside the telescope beam. The conversion efficiency of flux density to antenna temperature (${\\rm T_{A}^{*}}$), 2.8 Jy ${\\rm K^{-1}}$, was used for the calculations. The observational results show the data combined on each day and the entire days during the observation period.  POG reported that the water line detections of Titan and Hyperion were $\\sim$30 mK ($3.8 \\sigma$ for 8 hour integration) and $\\sim$50 mK ($4.0 \\sigma$ for 6 hours integration), respectively, which correspond to 300 mJy and 500 mJy, respectively. They found that these emissions were seen on almost the same satellite positions for both objects, and based on these results, they inferred a common mechanism of the emission such as Saturnian magnetosphere bow shock. On 16-18 and 28 May 2009, Titan and Hyperion were near conjunction so that we could observe them simultaneously in the telescope beam. During 16-18 May, especially, these objects were at almost the same position. Therefore, if the maser emissions had been caused by the common mechanism, we would have detected them simultaneously for these two objects. However, our observations this time could not detect the signals of the emissions, and all the data combined also showed no symptom. We estimated typical values of the $3\\sigma$ upper limit for each daily data. We also obtained $3\\sigma$ upper limits during the observations using all the daily data combined. The typical(all) $3\\sigma$ upper limits were 200(60) mJy for Titan and 240(80) mJy for Hyperion, which means that we could certainly detect the line emissions if the levels of flux densities were as high as those of POG.  Although POG did not report the water maser emission on Iapetus, we monitored this satellite during our observation period aiming this object at the OFF-point of the position switching observations. We could not find any signal stronger than $\\sim$210(75) mJy ($3\\sigma$ upper limit) .  The ${\\rm H_{2}O}$ ice and vapor plume on Enceladus has been reported by the Cassini Ultraviolet Imaging Spectrometer (Hansen {\\it et al.} 2006 and Hansen {\\it et al.} 2008), and the hypothesis that there exists liquid water in the crust has been proposed (Porco {\\it et al.} 2006). As for the water maser emission, the mass ratio of water vapor to ice in the plume is important because a majority of water molecules must be in the form of vapor in order to have maser emission originated from the Enceladus plume. A recent theoretical study indicates that the plume would be dominated by vapor from the thermodynamics perspective (Kieffer {\\it et al.} 2009). The maser line intensity reported by POG was 500 mJy (4.2$\\sigma$), and they estimated  the column density of water vapor from the observed maser intensity, which agrees with the column density of the water vapor plume observed in UV ($n{\\rm =1.5 \\times 10^{16}~cm^{-2}}$: Hansen {\\it et al.} 2006). Nevertheless, we did not find  appreciable maser emission in our data. The $3\\sigma$ upper limits were 780(540) mJy. Our data were worse compared with those of POG because Enceladus was located around the edge of the telescope beam much of the observational time. However, these results might indicate that the flux of the plume is varying, or the plume is sporadic; like a geyser.  POG reported the most certain detections for Atlas. The averaged spectrum showed a peak with 32 mK and S/N=7.0. From the satellite positions, they found that the maser emissions occurred on the trailing side which was several thousand km away from Atlas, and suggested that the disturbance of the Atlas's motion had caused the emission in the edge regions of Saturnian rings A and F. We attempted to verify this subject in our data. We divided the positions of Atlas into 3 parts; $i.e.$, position 1 and 3: positions passing before the geometrical edges of the Saturn's rings seen from the Earth, and Position 2: those after the rings (see the positions of Atlas in Figure 2). For each position, we combined the data, and furthermore, we tested the data on 19 May because Atlas passed by the geometrical edges of the rings ($\\theta=90^{\\circ}$) during the observations. However, we could not have positive results for each case. All the combined data did not show any prominent feature of the emission. The 3$\\sigma$ detection limits for each data set were in the range of 430--520(400) mJy.  We also checked several inner satellites which have a diameter larger than $\\sim$10 km;  Mimas, Janus, Epimeteus, Prometheus, Pandora and Pan. For all the satellites, the acquired data did not show indicative of the maser emission. The typical value of $3\\sigma$ upper limits were 500(210) mJy.  We had an opportunity to observe the Saturnian water maser line at the time of the ring's disappearance; that is, we see the Saturn ring almost in the edge-on view, and we obtained the data which had comparable or even better sensitivities than those by POG for almost of the satellites, though our observations were made for a limited period in 2009 May.  If maser emission is stationary in intensity at the levels as those reported by POG, we could detect them. However the results were negative for all of the observed Saturnian satellites. From these results, we conclude that the water maser  in the Saturnian system may be sporadic in nature and it is strongly restricted to the time and position of satellites. We have to monitor the satellites for  longer periods and to obtain statistical results. These studies would be useful to figure out the water maser emission reported in the Saturnian system. In addition, we should also perform monitoring observations for other icy bodies; Jovian satellites, comets, outer asteroids and Kuiper belt objects would be inside the scope. As suggested by Cosmovici {\\it et al.} (1996), the maser emission may be induced by catastrophic events. Such events like a disruption or eruption can occasionally be found among the solar system objects; ($e.g.$ C/1999 S4 Linear: disruption and dissipation, 29P/Schwassmann-Wachmann: outburst, 7968 Elst-Pizarro: impact or cometary activity (Toth 2000 and Hsieh et al. 2004), etc.). We should not miss events which will occasionally happen and carry out the observations to accumulate the data."
    },
    "1005/1005.3829_arXiv.txt": {
        "abstract": "We report the first X-ray detection of \\lya\\ emitters at redshift z$\\sim$4.5. One source (J033127.2-274247) is detected in the Extended Chandra Deep Field South (ECDF-S) X-ray data, and has been spectroscopically confirmed as a $z = 4.48$ quasar with $L_X = 4.2\\times 10^{44} \\ergsec$. The single detection gives a \\lya\\ quasar density of $\\sim$ 2.7$^{+6.2}_{-2.2}\\times$10$^{-6}$ Mpc$^{-3}$, consistent with the X-ray luminosity function of quasars.  Another 22 \\lya\\ emitters (LAEs) in the central Chandra Deep Field South (CDF-S) region are not detected individually, but their coadded counts yields a S/N=2.4 (p=99.83\\%) detection at soft band, % with an effective exposure time of $\\sim 36$ Ms. Further analysis of the equivalent width (EW) distribution shows that all the signal comes from 12 LAE candidates with EW$_{rest} <$ 400 \\AA, and 2 of them contribute about half of the signal. From follow-up spectroscopic observations, we find that one of the two is a low-redshift emission line galaxy, and the other is a Lyman break galaxy at z = 4.4 with little or no Ly$\\alpha$ emission. Excluding these two and combined with ECDF-S data, we derive a 3-$\\sigma$ upper limit on the average X--ray flux of $F_{0.5-2.0 keV}$ $<$ 1.6 $\\times 10^{-18}$ \\fluxunit,  which corresponds to an average luminosity of $\\langle L_{0.5-2 keV} \\rangle$ $<$ 2.4 $\\times 10^{42}$ ergs $s^{-1}$ for z $\\sim$ 4.5 \\lya\\ emitters.  If the average X-ray emission is due to star formation, it corresponds to a star-formation rate (SFR) of $<$ 180--530 M$_\\sun$ yr$^{-1}$. We use this SFR$_X$ as an upper limit of the unobscured SFR to constrain the escape fraction of \\lya\\ photons, and find a lower limit of f$_{esc,Ly\\alpha}$ $>$ 3--10\\%. However, our upper limit on the SFR$_X$ is $\\sim$7 times larger than the upper limit on SFR$_X$ on z$\\sim$ 3.1 LAEs in the same field, and at least 30 times higher than the SFR estimated from \\lya\\ emission. From the average X-ray to \\lya\\ line ratio, we estimate that fewer than 3.2\\% (6.3\\%) of our LAEs could be high redshift type 1 (type 2) AGNs, and those hidden AGNs likely show low rest frame equivalent widths. ",
        "introduction": "The LAE candidates were selected with narrowband imaging of the GOODS CDF-S (RA 03:31:54.02, Dec -27:48:31.5, J2000)  at the Blanco 4m telescope at Cerro Tololo InterAmerican Observatory (CTIO) with the MOSAIC II camera. Three 80 \\AA\\ wide narrowband filters (NB656, NB665 and NB673) were utilized to obtain deep narrowband images (Finkelstein et al. 2008, 2009c). The LAE candidates are  selected based on a 5 $\\sigma$ detection in the narrowband, a 4 $\\sigma$ significant narrowband flux excess over the broad band continuum image (here, an R band image from the ESO Imaging Survey [EIS], Arnouts et al. 2001), a factor of 2 ratio of narrowband flux to broadband flux density, and no more than 2 $\\sigma$ significant flux in the EIS-B band. Candidates with GOODS B-band coverage were further examined in the GOODS B-band image, and those with significant B-band detections were excluded. These conditions are satisfied by 113 LAE candidates with the \\textit{Chandra} CDF-S and ECDF-S coverage, including 4 in the NB656 filter\\footnote{The NB656 data was much shallower than the other two bands, thus the galaxies were selected in a different way (see Finkelstein et al. 2008) - we search for the NB656 candidates from the positions of galaxies which were detected in GOODS V-band but not in GOODS B-band.  Thus, we were only able to select galaxies over the GOODS region, which is why only four objects were selected. The other two catalogs consist of all selected candidates over the overlap region between the MOSAIC image and the ESO Imaging Survey, which consists of a much larger area. } (% Finkelstein et al. 2008),  39  in NB665, and 81 in NB673 (% including 11 that were detected in both NB665 and NB673). The equivalent widths (EWs) of our LAEs were calculated from our narrowband and EIS-R broadband data. Finkelstein et al (2008, 2009c) have previously studied the 14 objects from this sample that lie within the GOODS \\textit{HST} field. For these sources, we choose the deeper GOODS V-band to calculate the EWs.  The 2 Ms  \\textit{Chandra X--Ray Observatory} ACIS (Advanced CCD Imaging Spectrometer) exposure of the CDF-S is composed of 23 individual ACIS-I observations. We downloaded the raw data from the \\textit{Chandra} public archive and reduced the data using the \\textit{Chandra}  Interactive Analysis of Observations software version 4.0 (CIAO4.0). Each observation was filtered to include only standard \\textit{ASCA} event grades 0, 2, 3, 4, 6. Cosmic ray afterglows, ACIS hot pixels, and bad pixels were removed, along with all data taken during high background time intervals. All exposures were then added to produce a combined  event file with a net exposure of 1.9 Ms. The  \\textit{Chandra} exposure of the ECDF-S is composed of 9 individual ACIS-I observations obtained in 2004, covering $\\sim$ 0.3 deg$^2$ with four pointings. We reprocessed the X-ray raw data of the four pointings separately. The averaged net exposure per pointing at ECDF-S was 238 ks. The aspect offset\\footnote{http://cxc.harvard.edu/cal/ASPECT/fix\\_offset/fix\\_offset.cgi} of both CDF-S and ECDF-S data was examined and no offset above 0.1\\arcsec\\ was found in either field. We used the published X-ray source catalogs of the 2-Ms CDF-S (Luo et al. 2008) and the 240-ks ECDF-S (Lehmer et al. 2005) in the following source-match and source-mask processes. ",
        "conclusions": "Our work shows that X-ray observation is an effective method to identify AGN, as well as foreground objects in LAE samples. One X-ray detected LAE is spectroscopically confirmed as a type 1 quasar at z = 4.5.   A stack of 22 other LAEs in the CDF-S field yields a marginal detection.   However, two of these 22 sources contribute about half of the stacked X-ray signal, and these two were found to be a foreground interloper and a LBG at z=4.4 without strong \\lya\\ emission. The mean flux of the remaining 20 sources, while positive, is not significantly different from zero. Including the ECDFS data,  we obtain a $3\\sigma$ upper limit on the average X-ray luminosity of  2.4 $\\times 10^{42}$ erg s$^{-1}$. Compared to their average \\lya\\ luminosity, we estimate that that fewer than 3.2\\% (6.3\\%) of our LAEs could be high redshift type 1 (type 2) AGNs, and those hidden AGNs might show low EW$_{rest}$. Using the relationship of X-ray emission and star-forming activity from low redshift star-forming galaxies, we obtained an upper limit on the unobscured SFR of SFR$<$ 180-530 M$_\\sun$ yr$^{-1}$. Compared to the SFR estimated from their average \\lya\\ luminosity, we find a lower limit on the escape fraction of \\lya\\ photons, f$_{esc,Ly\\alpha} >$ 3-10\\%. Doubling the depth of CDF-S X-ray observations is planned in 2010 and 2011 (see \\textit{Chandra} Electronic Bulletin 89). This will strengthen the power of X-ray diagnostics of LAEs, especially for revealing their unobscured SFR, for the new discovery of \\lya\\ selected quasars and weak AGN, and for excluding the low-redshift contamination."
    },
    "1005/1005.2191_arXiv.txt": {
        "abstract": "We present new moderate-resolution, far-ultraviolet spectra from the {\\it Hubble Space Telescope}/Cosmic Origins Spectrograph (HST/COS) of the BL\\,Lac object 1ES\\,1553$+$113 covering the wavelength range $\\rm 1135~\\AA<\\lambda<1795~\\AA$.  The data show a smooth continuum with a wealth of narrow ($b<100$ \\kms) absorption features arising in the interstellar medium (ISM) and intergalactic medium (IGM).  These features include 41 \\Lya\\ absorbers at $0<z_{\\rm abs}<0.43$, fourteen of which are detected in multiple Lyman lines and six of which show absorption in one or more metal lines.  We analyze a metal-rich triplet ($\\Delta cz\\sim1000$ \\kms) of \\Lya\\ absorbers at $z_{\\rm abs}\\approx0.188$ in which \\OVI, \\NV, and \\CIII\\ absorption is detected.  Silicon ions (\\SiIII, \\SiIV) are not detected to fairly strong upper limits, and we use the measured \\SiIII/\\CIII\\ upper limit to derive an abundance limit $\\rm(C/Si)\\ge4\\,(C/Si)_\\sun$ for the strongest component of the absorber complex.  Galaxy redshift surveys show a number of massive galaxies at approximately the same redshift as this absorption complex, suggesting that it arises in a large-scale galaxy filament.   As one of the brightest extragalactic X-ray and $\\gamma$-ray sources, 1ES\\,1553$+$113 is of great interest to the high-energy astrophysics community.  With no intrinsic emission or absorption features, 1ES\\,1553$+$113 has no direct redshift determination.  We use intervening \\Lya\\ absorbers to place a direct limit on the redshift: $z_{\\rm em}>0.395$ based on a confirmed \\Lya$+$\\OVI\\ absorber and $z_{\\rm em}>0.433$ based on a single-line detection of \\Lya.  The current COS data are only sensitive to \\Lya\\ absorbers at $z<0.47$, but we present statistical arguments that $z_{\\rm em}\\la0.58$ (at a $1\\sigma$ confidence limit) based on the non-detection of any \\Lyb\\ absorbers at $z>0.4$. ",
        "introduction": "The current interpretation of BL\\,Lac objects \\citep{Ghisellini85} is that they are active galactic nuclei (AGN) with a strongly relativistic jet pointed toward our line of sight.  As such, any line emission or accretion disk features seen in most other types of AGN could be masked by the bright jet if present in BL Lac objects.  Their spectra usually show a featureless power-law continuum extending from radio to X-ray wavelengths.  This spectral characteristic makes BL\\,Lac objects ideal for observing intervening absorption features arising in the interstellar medium (ISM) and intergalactic medium (IGM).  Since their continuum is easily defined, they make excellent targets for studying weak metal-line systems and low-contrast, highly thermally broadened \\HI\\ absorbers \\citep[e.g.,][]{Richter04,Lehner07,Danforth10}.  The BL\\,Lac object 1ES\\,1553$+$113 shows the characteristic featureless power-law spectrum and is one of the brightest known sources of extragalactic high-energy radiation from X-rays up to VHE (TeV) photons \\citep{CostamanteGhisellini02}.  However, the featureless spectrum makes it difficult to determine the redshift of the object and hence its luminosity.  Indirect methods have given a wide range of limits for the redshift of 1ES\\,1553$+$113; the nondetection of a host galaxy gave limits from $z_{\\rm em}>0.09$ to $z_{\\rm em}>0.78$ \\citep{HutchingsNeff92,Scarpa00,Urry00,Carangelo03,Sbarufatti06,Treves07}.  The shape of the $\\gamma$-ray spectrum observed by the {\\it Fermi Observatory} and ground-based VHE detectors (HESS, MAGIC) constrains the redshift to values from $z_{\\rm em}<0.4$ to $z_{\\rm em}<0.8$ \\citep{Ahronian06,Albert07,MazinGoebel07,Abdo10} based on assumptions about the intrinsic spectral energy distribution (SED) and pair-production interactions with the cosmic infrared background.  The only direct redshift determination \\citep[$z_{\\rm em}=0.36$;][]{MillerGreen83} was based on a spurious feature in low-resolution UV spectra from the {\\it International Ultraviolet Explorer} (IUE).  The detection was later retracted \\citep{FalomoTreves90}, but the erroneous redshift value lives on.  1ES\\,1553$+$113 is of interest as a bright background continuum source for detecting intergalactic absorption along the sight line.  Bright X-ray sources are especially valuable for potentially detecting the long-predicted \\OVII\\ and \\OVIII\\ tracers \\citep{Bregman07} of intergalactic gas at $T=10^6-10^7$~K.  Even for a bright X-ray source, the required integration times would be very long.  However, a sufficiently long IGM pathlength provided by a bright high-$z$ target would make the required observing time investment more attractive.  In this paper, we present the first medium-resolution far-UV spectroscopic observations of 1ES\\,1553$+$113 including {\\it Hubble Space Telescope}/Cosmic Origins Spectrograph \\citep[HST/COS][]{Green10,Osterman10} observations ($\\lambda=1135-1795$~\\AA) as well as archival data at $905-1187$~\\AA\\ from the {\\it Far Ultraviolet Spectroscopic Explorer} \\citep[\\FUSE;][]{Moos00,Sahnow00}.  We confirm the featureless power-law nature of the spectrum over this wavelength range.  Absorption is seen in 42 intervening systems including 41 \\Lya\\ absorbers and six metal-line systems.  The frequency of IGM absorbers is consistent with larger surveys using \\FUSE\\ and HST/STIS data (Danforth \\& Shull 2005, 2008; hereafter DS08), and the systems are spread across the entire redshift range covered by the combined COS/\\FUSE\\ dataset ($z\\la0.47$).  The observations and data reduction techniques are discussed in \\S2, and we present a preliminary catalog of absorption lines in \\S3.  Our conclusions are presented in \\S4. ",
        "conclusions": "The BL\\,Lac object 1ES\\,1553$+$113 is one of the brightest objects in the sky in $\\gamma$-rays, as well as being a notable UV and X-ray source.  However, the AGN emission is that of a relativistic jet aligned closely with our line of sight and, like most such objects, has no intrinsic emission or absorption features at any wavelength.  This featureless, power-law continuum is ideal for measuring intervening IGM features that are weak and broad, such as thermally-broadened \\Lya\\ systems.  However, the lack of intrinsic features makes constraining the redshift of the object difficult.  We present unprecedented high-quality far-UV HST/COS and \\FUSE\\ spectra of the BL\\,Lac object 1ES\\,1553$+$113 at spectral resolution 15-20 \\kms.  These data show 42 intervening IGM absorbers, 41 of which are detected in \\Lya, and 15 in \\Lyb\\ and/or metal lines.  The richest absorption system in the line of sight is a trio of \\Lya\\ absorbers at $z\\approx0.188$ covering $\\sim1000$ \\kms\\ of velocity space.  Several metal ions are also detected in these systems, including \\OVI, \\NV, and \\CIII.  However, neither \\SiIV\\ nor \\SiIII\\ is detected in any of the systems.  The \\CIII/\\SiIII\\ ratio implies a (C/Si) abundance at least four times the solar value, while a high \\NV/\\OVI\\ value suggests an overabundance of N as well.  A detailed analysis of the physical conditions in this system can be found in \\citet{Yao10}.  The redshift of 1ES\\,1553$+$113 has never been determined directly, and the only limits placed on it come from indirect means such as the shape of the $\\gamma$-ray spectrum and the lack of an AGN host galaxy in deep optical images.  A strong \\Lya$+$\\OVI\\ absorber at $z=0.3951$ gives the first direct lower limit to the redshift of the object.  Two weaker \\Lya\\ absorbers at $z=0.4063$ and $z=0.4326$ give slightly higher estimates of the redshift, but these weak \\Lya\\ lines are not confirmed by additional line detections.  These lower limits are consistent with most previous measurements via optical non-detections of host galaxies and $\\gamma$-ray SED constraints.  \\citet{Abdo10} derive $z_{\\rm em}=0.75^{+0.04}_{-0.05}$ based on the latest {\\it Fermi} and TeV $\\gamma$-ray SED measurements, considerably higher than our intervening absorber upper limits.  COS far-UV spectra are not sensitive to \\Lya\\ absorbers at $z>0.47$, but the G160M grating has some sensitivity to intervening \\Lyb\\ and \\OVI\\ absorbers out to $z\\sim0.75$.  If the \\citet{Abdo10} redshift estimate were accurate, we would expect to find $\\sim8$ \\Lyb\\ absorbers at $0.47<z<0.75$.  We find no evidence for any higher redshift absorption systems.  There are only a few absorption features at $\\lambda>1500$~\\AA\\ with ambiguous line identifications that could potentially be \\Lyb\\ systems at $z>0.47$.  While these systems are individually suggestive, we find nowhere near the number of absorbers predicted statistically.  We conclude that the redshift of 1ES\\,1553$+$113 is not much higher than $z\\approx0.45$.  1ES\\,1553$+$113 is one of the brightest X-ray sources on the sky and has been suggested as a sight line that could be efficiently probed for WHIM absorption in \\OVII.  The combined \\OVI\\ column density in the three absorbers at $z\\sim0.19$ is $\\sim2\\times10^{14}\\rm~cm^{-2}$.  Spectrographs on modern X-ray observatories are sensitive to $\\log\\,N_{\\rm OVII}\\ga15.5$.  If the temperature of any of these \\OVI\\ systems is high enough, sufficiently long Chandra and/or XMM/Newton observations may reveal a \\OVII\\ counterpart to these \\OVI\\ absorbers that could constrain the long-sought X-ray WHIM \\citep{Bregman07}.  However, at the observed Li-like (\\OVI) oxygen column density, $\\log\\,N_{\\rm OVI}\\approx14.3$ in the trio of absorbers, the expected column densities of He-like (\\OVII) and H-like (\\OVIII) oxygen are probably just below the detectability levels of {\\it Chandra} and {\\it XMM}.  Recent analysis of stacked X-ray absorption data \\citep{Yao09} at the known IGM redshifts of \\OVI\\ absorbers finds no evidence for \\OVII\\ or \\OVIII\\ absorbers to a limit $N_{\\rm OVII}/N_{\\rm OVI}<10$.  Therefore, the $z\\approx0.19$ absorbers might have X-ray column densities $\\log\\,N_{\\rm OVII}\\leq15.3$, just below the limits of current X-ray observatories.  Finally, these observations showcase the powerful new tool available to astronomers for probing the low-redshift IGM.  COS is $10-20$ times more sensitive in the far-UV to point sources than previous instruments on HST.  An additional six orbits of COS observations are planned for 2010, which should improve the S/N of the combined dataset by a factor of $\\sim\\sqrt{3}$.  Improving the data quality will help confirm or refute some of the tentative line identifications from this paper and will undoubtedly uncover additional weak absorbers.  We will place further constraints on [C/Si] and [N/O] in the $z=0.188$ system, identify new broad, shallow \\Lya\\ absorbers, and investigate possible high-$z$ \\Lyb\\ systems with our new Cycle 18 observations.  \\medskip \\medskip It is our pleasure to acknowledge the many thousands of people who made the HST Servicing Mission 4 the huge success that it was.  We furthermore thank Steve Penton, St\\'ephane B\\'eland, and the other members of the COS ERO and GTO teams for their work on initial data calibration and verification.  C. D. wishes to acknowledge a fruitful discussion with members of the KIPAC consortium.  This work was supported by NASA grants NNX08AC146 and NAS5-98043 to the University of Colorado at Boulder."
    },
    "1005/1005.2969_arXiv.txt": {
        "abstract": "{ High-resolution far-infrared and sub-millimetre spectroscopy of water lines is an important tool to understand the physical and chemical properties of cometary atmospheres. We present  observations of several rotational ortho- and para-water transitions in comet \\garradd{} performed with HIFI on {\\it Herschel}. These observations have provided the first detection of the $2_{12}$--$1_{01}$ (1669 GHz) ortho and $1_{11}$--$0_{00}$ (1113 GHz) para transitions of water in a cometary spectrum. In addition, the ground-state transition $1_{10}$--$1_{01}$ at 557 GHz is detected and mapped.  By detecting several water lines quasi-simultaneously and mapping their emission we can constrain the excitation parameters in the coma. Synthetic line profiles are computed using excitation models which include excitation by collisions, solar infrared radiation, and radiation trapping. We obtain the gas kinetic temperature, constrain the electron density profile, and estimate the coma expansion velocity by analyzing the map and line shapes. We derive water production rates of $\\range \\times10^{28}\\ \\mathrm{s}^{-1}$ over the range $r_\\mathrm{h}=1.83$--$1.85$ AU.} ",
        "introduction": "Comets spend most of their lifetime in the outer Solar System and therefore have not undergone considerable thermal processing.  Line emission is useful to study the physical and chemical conditions of cometary atmospheres, and their relation to other bodies in the Solar System \\citep{2002EM&P...90..323B,2004come.book..391B}.  Water molecules in cometary atmospheres are excited due to collisions with other molecules and radiative pumping of the fundamental vibrational levels by the solar infrared flux. The \\transi{} ortho-water transition at 557 GHz is one of the strongest lines in cometary comae, but it cannot be detected directly from the ground due to absorption in the Earth's atmosphere \\citep{1987A&A...181..169B}. Water vapour production has been estimated previously from the ground through measurements of its photodissociation product, the OH radical \\citep[see e.g.][]{1982come.coll..433A}, and water high vibrational bands \\citep{2004come.book..391B}. The \\transi{} rotational transition of ortho-water at 557 GHz has been observed using heterodyne techniques by the Submillimeter Wave Astronomical Satellite (SWAS) \\citep{2000ApJ...539L.151N,2001Icar..154..345C}, and later with Odin \\citep{2003A&A...402L..55L,2007P&SS...55.1058B,2009A&A...501..359B}.  ESA's {\\it Herschel} Space Observatory was successfully launched on May 14, 2009 and entered a Lissajous orbit around the $L_2$ Lagrangian point \\citep{2010Herschel}. The Heterodyne Instrument for the Far-Infrared (HIFI) onboard {\\it Herschel} has continuous coverage in five frequency bands in the 480--1150 GHz range, and an additional dual frequency band covering the 1410--1910 GHz range that are not observable from the ground \\citep{2010HIFI}.  HIFI's submillimetre frequency coverage is of great importance to observe water vapour in Solar System objects such as cometary comae with unique sensitivity and the required high spectral resolution to resolve the emission lines.  Comet \\garradd{} is a long-period comet ($P = 1.9\\times10^5$) with a highly eccentric orbit ($e=0.99969$). It passed perihelion on June 23, 2009 at a distance of 1.7982 AU from the Sun and was observed with HIFI in July 2009 as part of the {\\it Herschel} guaranteed time key project ``Water and related chemistry in the Solar System'' \\citep{2009P&SS...57.1596H}.  In this letter, we describe the observations of water and the analysis of part of the data set. Water production rates derived from our radiative transfer models are presented. ",
        "conclusions": "\\label{sec:conclusions}  The {\\it Herschel} Space Observatory provides unique new capabilities for the detection of water in the Solar System.  HIFI's spectral range and high resolution allowed for the direct detection of several water lines almost simultaneously.  High spectral resolution is crucial to resolve the line shape and asymmetries due to self-absorption.  On 20-27 July 2009, comet \\garradd{} was observed with HIFI. The high-resolution spectra of HIFI allows us to detect for the first time several rotational water lines in cometary spectra.  A water production rate of $\\sim 2\\times10^{28}\\ \\mathrm{s}^{-1}$ was derived at heliocentric distance of 1.8 AU using radiative transfer numerical codes which include collisional effects and infrared fluorescence by solar radiation.  In future studies, HIFI will be able to detect water isotopes, and determine the D/H ratio in active comets."
    },
    "1005/1005.5252_arXiv.txt": {
        "abstract": "{Non-blazar AGN have been recently established as a class of gamma-ray sources. M87, a nearby representative of this class, show fast TeV variability on timescales of a few days.} {We suggest a scenario of flare gamma-ray emission in non-blazar AGN based on a red giant interacting with the jet at the base.} {We solve the hydrodynamical equations that describe the evolution of the envelope of a red giant blown by the impact of the jet.} {If the red giant is at least slightly tidally disrupted by the supermassive black hole, enough stellar material will be blown by the jet, expanding quickly until a significant part of the jet is shocked. This process can render suitable conditions for energy dissipation and proton acceleration, which could explain the detected day-scale TeV flares from M87 via proton-proton collisions. Since the produced radiation would be unbeamed, such an events should be mostly detected from non-blazar AGN. They may be frequent phenomena, detectable in the GeV-TeV range even up to distances of $\\sim 1$~Gpc for the most powerful jets. The counterparts at lower energies are expected to be not too bright.} {M87, and nearby non-blazar AGN in general, can be fast variable sources of gamma-rays through red giant/jet interactions.} ",
        "introduction": "\\label{intro}  Active galactic nuclei (AGN)  are believed to be powered by an accreting supermassive black hole (SMBH) in the center of a galaxy, a significant fraction of AGN show powerful jets, supersonic relativistic flows,  on small (sub parsec) and large (multi hundred kpc) scales. \\citep[e.g.][]{bbr84}.  These AGN are characterized by nonthermal emission extending from radio to high energy gamma-rays. This radiation comes from an accretion disc and from two relativistic jets that are launched close to the SMBH in two opposite directions. The emission associated to the accretion process can be generated by thermal plasma in the form of an optically-thick disc under efficient cooling \\citep[e.g.,][]{s72,ss73}, or as an optically-thin corona \\citep[e.g.,][]{bkb77,lt79}. The emission from the jets is non thermal and comes from a population of relativistic particles accelerated for instance in strong shocks, although other scenarios are possible as well \\citep[see, e.g.,][]{recon,na07,rbd07,ra08}. This non-thermal emission is thought to be produced through synchrotron and inverse Compton (IC) processes \\citep[e.g.][]{gmt85}, although hadronic models have been also considered in the past \\citep[e.g.,][]{m93,ah00,mp01,a02}.  The existence of a stellar clustering in the central regions of AGN, possibly down to very small distances from the central SMBH \\citep[e.g.][]{pens88}, implies that the interaction between a star and the jet should eventually happen.  The gamma-ray production due to the interaction between an obstacle and an AGN jet has been studied in a number of works. For instance, \\cite{dl97}  suggested the high-energy radiation produced by a beam of relativistic protons impacting with a cloud of the broad-line region (BLR). The gamma-ray emission from one or many clouds from the BLR interacting with a hydrodynamical jet recently has bin analyzed by \\cite{abr10}. The radiation from the interaction between a massive star and an AGN jet was studied in \\cite{bp97}. Namely they suggested that the jet interacts with stellar winds of massive stars, in their model they assume that the source of gamma-rays is moving with a relativistic speed, therefore  the radiation is  Doppler boosted. The main radiation mechanism in this scenario is related to the development of the pair cascade in the field of the radiation of the  massive star.  In this work, we study the interaction of a red giant (RG) star with the base of the jet in AGN and their observable consequences in gamma rays. We focus here on the case of M87, a nearby non-blazar AGN that presents very high-energy recurrent activity with variability timescales of few days \\citep{ah06,mag,vvhm09,ver10}. In the framework presented here, the jet impacts the RG envelope, already partially tidaly disrupted by the gravitational field of the central SMBH. The RG envelope is blown up, forming a cloud of gas accelerated and heated by jet pressure. The jet base is likely strongly magnetized \\citep[e.g.,][]{kbvk07,bk08}. The jet flow affected by the impact with the RG envelope can be a suitable region for particle acceleration, and a significant fraction of the involved magnetic and kinetic energy of the jet can be transferred to protons and electrons. Although electrons may not able to reach TeV emitting energies because of the expected large magnetic fields, protons would not suffer from this constraint. These protons could reach the star blown material, and optically-thick proton-proton ($pp$) interactions could lead to significant gamma-ray production in the early stages of the cloud expansion. Unlike in \\cite{bp97}, we deal with solar-mass-type stars instead of the more rare high-mass stars, study the RG atmosphere-jet interaction, and follow the hydrodynamical evolution of the cloud. Finally, we do not introduce any beaming factor to the radiation, since in our scenario most of the emission is produced when the cloud has not been significantly accelerated, Doppler boosting being therefore negligible. ",
        "conclusions": "The total jet luminosity can be inferred from observations using Eq.~(\\ref{egcn}): \\begin{equation} L_{\\rm j}=8\\times 10^{44}\\,L_{\\gamma,41}^{5/8} M_{\\rm BH,9}^{2/3}\\theta_{-1}\\, \\eta_{-1}^{-5/8}t_{\\rm v,5}^{-1/8}M_{\\rm RG}^{-1/6}\\,\\mbox{erg s}^{-1}. \\label{etj} \\end{equation} This formula weakly depends on the observables, being almost insensitive to $M_{\\rm RG}$, on the other hand hard to estimate. This provides therefore a quite robust estimate of the jet luminosity with $\\eta$ as most unknown parameter. Actually, if $L_{\\rm j}$ were known, then $\\eta$ could be also estimated.  For the most powerful jets, $L_\\gamma$ would be limited by the jet size becoming $L_\\gamma=\\chi\\eta L_{\\rm j}$. Taking for instance $L_{\\rm j}\\sim 10^{47}$~erg~s$^{-1}$, $L_\\gamma$ could be as high as  $\\approx 2\\times 10^{45}\\,\\eta_{-1}$~erg~s$^{-1}$.  An improvement of a factor of several in the VHE sensitivity (e.g. through the forthcoming the Cherenkov Telescope Array -CTA-) would test our gamma-ray predictions for the whole RG-jet interaction process, including the early cloud expansion phase, allowing for a detailed study of the involved (magneto)hydrodynamics, particle acceleration, and radiation.  We remark that, if a detectable gamma-ray flare with a duration of few days were to be produced in M87, in particular through $pp$ interactions, the cloud should have a mass of $\\sim 10^{28}$~g. Such a massive cloud cannot acquire a large speed in the jet direction at the times when $pp$ collisions are an efficient gamma-ray emitting mechanism, and therefore the emission will not suffer significant Doppler boosting. In the case of a lighter cloud, large Lorentz factors can be achieved, but then $pp$ interactions will be inefficient producing gamma-rays, the probability to detect a flare lower due to beaming, and the duration of the event shorter than observed because of faster expansion and beaming.  {  Coming back to the question of cloud mass, we note that to extract a cloud with a mass $>10^{28}$ g, a more powerful jet than in M87, for similar jet-RG interaction conditions, would be required (see Eq.(\\ref{mb})). }  An important question is whether there are enough RGs in M87 at the relevant jet scales. The model presented here would require few interactions per year to explain the observations in M87. Since the typical duration of the RG-jet interaction is of about 3--4~days, the RG filling factor should be $\\Upsilon\\sim 4/365\\approx 10^{-2}$. With a jet volume at the relevant scales of $\\sim \\pi\\theta^2 z_{\\rm T}^3/3$, the density of RGs in the region should be $\\sim \\Upsilon/V\\sim 2\\times 10^6$~pc$^{-3}$ for M87. Unfortunately, no direct information is available on the density of stars in the vicinity of the SMBH in M87. The stellar mass in a sphere with a radius of 80~pc is estimated in $2\\times 10^8 M_{\\odot}$ \\citep[e.g.][]{gt09}, and these observational data should be extrapolated four orders of magnitude down to $\\sim 0.01$~pc. Thus, depending on the assumed extrapolation law, the number of RGs in the vicinity of the SMBH may or may not be enough. It is worth noting that a dense stellar cluster near the SMBH could be behind the broad-line region in AGN as produced by the blown-up atmosphere of red dwarfs, which would imply the presence of numerous RGs in the center of AGN \\citep{pens88}. In addition, studies of the possible stellar density profiles in the vicinity of the SMBH in AGN \\citep{bkck82,mcd91} show that densities like the required one ($\\sim  2\\times 10^6$~pc$^{-3}$) could be achieved. The observation of VHE flares could be already an indication that enough RGs are present near the SMBH in M87.  Interestingly, RG/jet interactions are expected to be transient phenomena. At higher jet heights, although many RGs could be simultaneously present in the jet rendering rather continuous emission, the much more diluted jet would not remove a significant amount of material from the star and the effective cross section of the interaction would be just $R_{\\rm RG}$, yielding a low energy budget for such a multiple interaction events.  The scenario presented here, adopted to explain the day VHE flares observed from M87, could also be relevant in other non-blazar AGN. For blazar sources the beamed emission would overcome the RG-jet interaction, expected to be weakly beamed due to moderate $v_{\\rm z}$-values. For instance, the closest AGN, the radio galaxy Cen~A, at $\\sim 3.8$~Mpc distance \\citep{r04}, could also show detectable flare like emission. At present, persistent faint VHE emission has been detected \\citep{ah09} with $L_{\\gamma}=2.6\\times 10^{39}\\,\\mbox{erg~s}^{-1}$. Accounting for the black hole mass of this AGN, $M_{\\rm BH}=5.5\\times 10^7\\,M_{\\odot}$, taking the observed VHE luminosity as a reference, and assuming $t_{\\rm v}\\sim 1$~day, one derives implementing Eq.~\\ref{etj} a jet luminosity $L_{\\rm j}=1.2\\times 10^{42}\\,\\mbox{erg~s}^{-1}$, a rather modest value. Therefore, it cannot be excluded that RG-jet interactions may contribute to the VHE radiation detected from Cen~A, or that transient activity due to RG-jet interactions may be observed from this source. Another case, the radio galaxy NGC1275, at a distance of 73~Mpc \\citep{hwb09}, shows variable behaviour in GeV \\citep{fer09}. The GeV luminosity is of about $2\\times 10^{43}$~erg~s$^{-1}$. Using Eq.~(\\ref{etj}), we can estimate the power of the jet as $5\\times 10^{44}$~erg~s$^{-1}$ adopting an $M_{BH}=10^8\\,M_\\odot$. In the case of NGC~1275 the shocked cloud would be optically thick at the luminosity peak, implying significant attenuation of the TeV emission through photon-photon absorption with a cutoff around 50~GeV.  At farther distances, the strong jet luminosity dependence $L_{\\gamma}\\propto L_{\\rm j}^{1.6}$ implies that FR~II sources with say $L_{\\rm j}\\sim 10^{46}$~erg~s$^{-1}$ may be still detectable up to distances of $\\sim 0.5$~Gpc (internal absorption should be included in these cases; see \\citealt{akc08}). Also, the luminosity in the range 0.1--100~GeV would also be significant unless there is a strong low-energy cutoff in the proton spectrum. Therefore, Fermi may detect day-long GeV flares originated due to RG-jet interactions from FR~II galaxies up to distances of few 100~Mpc. Summarizing, GeV and TeV instrumentation can potentially detect a number of RG-jet interactions per year taking place in nearby FR~II and very nearby FR~I galaxies, with the most powerful events being detectable up to 1~Gpc."
    },
    "1005/1005.0641_arXiv.txt": {
        "abstract": "{We introduce two new methods that are designed to improve the realism and utility of large, active region-scale 3D MHD models of the solar atmosphere.  We apply these methods to RADMHD, a code capable of modeling the Sun's upper convection zone, photosphere, chromosphere, transition region, and corona within a single computational volume.  We first present a way to approximate the physics of optically-thick radiative transfer without having to take the computationally expensive step of solving the radiative transfer equation in detail.  We then briefly describe a rudimentary assimilative technique that allows a time series of vector magnetograms to be directly incorporated into the MHD system. ",
        "introduction": "In this paper, we briefly summarize our efforts to improve our models of quiet Sun and active region magnetic fields in computational domains that include the upper convection zone, photosphere, chromosphere, transition region and low corona within a single computational domain. Our goal is similar to that presented in \\citet{Abbett07} --- that is, to develop the techniques necessary to efficiently simulate the spatially and temporally disparate convection zone-to-corona interface over spatial scales sufficiently large to accommodate at least one active region.  The advantage of this type of single-domain modeling is clear. For example, evolving a turbulent convection zone and corona simultaneously in a physically self-consistent way allows for the quantitative study of important physical processes such as flux emergence, submergence and cancellation; the transport of magnetic free energy and helicity into the solar atmosphere; the generation of magnetic fields via a convective dynamo; and the physics of coronal heating.  However, this approach is challenging.  The computational domain is highly stratified --- average thermodynamic quantities change by many orders of magnitude as the domain transitions from a relatively cool, turbulent regime below the visible surface, to a hot, magnetically-dominated and shock-dominated regime high in the model atmosphere.  In addition, the low atmosphere is where the radiation field transitions from being optically thick to optically thin.  The chromosphere itself presents an additional challenge, since the radiation field is often decoupled from the thermal pool, particularly in some of the strongest, most energetically important transitions.  There are a number of ways to model the energetics of the convection zone-to-corona system, ranging from approximate, parameterized descriptions of the thermodynamics (see e.g., \\citealt{Fan09,Hood09}), to highly realistic treatments of radiative transfer (see e.g., \\citealt{Martinez09a, Martinez09b}). Since our objective is to model the coupled system over large spatial scales, our goal is to find the most efficient treatment of the energetics possible that still provides a physically meaningful representation of the dynamic connection between the convection zone and corona.  In order to describe the thermodynamics of the corona, a model should include the effects of electron thermal conduction along magnetic field lines and radiative cooling in the optically-thin limit. In addition, some physics-based or empirically-based source of coronal heating must be present if the model corona is to remain hot.  In the convective interior well below the visible surface, radiative cooling can be treated in the diffusion limit. The trick is, how best to describe the effects of optically-thick radiative transfer in the region of the model atmosphere that lies between these two extremes.  The most satisfying approach would be to couple the LTE transfer equation (or non-LTE population and transfer equations) to the MHD system to obtain cooling rates and intensities that could be compared directly to observations. Unfortunately, for large active region or global-scale problems, the computational expense of these techniques remains prohibitive.  In \\citet{Abbett07} we tried the opposite approach --- ignore the transfer equation altogether, and develop an artificial, fully parameterized means of approximating surface cooling (in this case, we employed a modified form of Newton cooling).  This worked relatively well, provided we carefully calibrated the adjustable parameters to match the average sub-surface stratification of previous, more realistic simulations of magneto-convection where the LTE transfer equation was solved in detail \\citep{Bercik02}.  Of course, the principle drawback of this approach is that it is ultimately \\emph{ad hoc} and unphysical, and requires other, more realistic simulations as a basis for calibration in order to get meaningful results.  We therefore have developed an approximation that is based on the macroscopic radiative transfer equation, and have incorporated this new treatment into our 3D MHD model, RADMHD.  We describe this new method in Section 2.  While it is important to treat the energetics of the system in a physically meaningful way, it is also important to remember that the utility of a given simulation ultimately depends on the statement of the problem.  For an MHD simulation, this boils down to one's choice of initial states and boundary conditions.  It is of great benefit, for example, to pose a simple, well-defined problem, and set up a numerical experiment that can shed light on what is believed to be the relevant physical processes in an otherwise complex system. For example, important progress has been made in understanding the physics of magnetic flux emergence by studying how idealized twisted flux ropes emerge through highly-stratified model atmospheres (see e.g., \\citealt{Cheung07,Fan04}).  Yet the observed evolution of the photospheric magnetic field is often far more complex, particularly in and around CME and flare producing active regions.  It is very difficult to set up a simple magnetic and energetic configuration that can initialize a simulation that will faithfully mimic the observed evolution of a real active region.  It is desirable to do so, however, since we wish to quantitatively understand the physical mechanisms of energy storage and release, and the transport of magnetic energy and helicity between the convective interior and corona.  To make progress, we could take a cue from meteorologists, and investigate a means to incorporate observational data directly into MHD models.  This is not at all straightforward for solar models however, since data is obtained entirely through remote sensing, and not \\emph{in situ}.  To address this challenge, we have developed a simple, rudimentary means of assimilating a time series of vector magnetograms into an MHD model of the photosphere-to-corona system.  We briefly summarize this technique in Section 3, and apply it to the specific problem of finding a 3D magnetic field that is as force-free as possible given a single measurement of the vector magnetic field at the photosphere. ",
        "conclusions": "We have developed a rudimentary means of assimilating a time series of vector magnetograms into the interior volume of an MHD model in a manner that is stable, and does not over-specify the problem. We are currently using this assimilative technique to incorporate a timeseries of vector magnetograms into a 120 Mm$\\,^3$ RADMHD model atmosphere that contains a model photosphere, chromosphere, transition region and corona.  The IVM data we are using is a four hour timeseries from NOAA AR 8210 --- a well-studied flare and CME-producing active region.  The simulations are in their preliminary stages, and we hope to report on this work in the near future.  In addition, we have presented a computationally efficient method of approximating optically-thick radiative cooling in our RADMHD quiet Sun models.   The treatment improves upon the method of \\citet{Abbett07}, while still retaining the efficiency necessary to allow for large, active region-scale, convection zone-to-corona computational domains.  Our simulations are progressing, and we are currently evaluating the efficacy and reliability of the new method.  We are optimistic that each of these methods will improve the realism and utility of our current suite of numerical models."
    },
    "1005/1005.1067_arXiv.txt": {
        "abstract": "Supermassive black holes are found at the centers of most galaxies and their inspiral is a natural outcome when galaxies merge. The inspiral of these systems is of utmost astrophysical importance as prodigious producers of gravitational waves and in their possible role in energetic electromagnetic events. We study such binary black hole coalescence under the influence of an external magnetic field produced by the expected circumbinary disk surrounding them. Solving the Einstein equations to describe the spacetime and using the force-free approach for the electromagnetic fields and the tenuous plasma, we present numerical evidence for possible jets driven by these systems. Extending the process described by Blandford and Znajek for a single spinning black hole, the picture that emerges suggests the electromagnetic field extracts energy from the orbiting black holes, which ultimately merge and settle into the standard Blandford-Znajek scenario. Emissions along dual and single jets would be expected that could be observable to large distances. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.3548_arXiv.txt": {
        "abstract": "Massive metal-poor stars might form massive stellar black holes (BHs), with mass $25\\le{}m_{\\rm BH}/{\\rm M}_{\\odot{}}\\le{}80$, via direct collapse. We derive the number of massive BHs (N$_{\\rm BH}$) that are expected to form per galaxy through this mechanism. Such massive BHs might power most of the observed ultra-luminous X-ray sources (ULXs). We  select a sample of  64 galaxies with X-ray coverage, measurements of the star formation rate (SFR) and of the metallicity. We find that N$_{\\rm BH}$ correlates with the number of observed ULXs per galaxy (N$_{\\rm ULX}$) in this sample. We discuss the dependence of our model on the SFR and on the metallicity. The SFR is found to be crucial, consistently with previous studies. The metallicity plays a role in  our model, since a lower metallicity enhances the formation of massive BHs.  Consistently with our model, the data indicate that there might be an anticorrelation between N$_{\\rm ULX}$, normalized to the SFR, and the metallicity. A larger and more homogeneous sample of metallicity measurements is required, in order to confirm our results. ",
        "introduction": "Ultra-luminous X-ray sources (ULXs, see Mushotzky 2004 for a review, and references therein) are defined as  non-nuclear point-like sources with isotropic X-ray luminosity $L_{\\rm X}\\gtrsim{}10^{39}$ erg s$^{-1}$. The mechanism that powers the ULXs is still unknown, although various scenarios have been proposed. ULXs could be associated with high-mass X-ray binaries (HMXBs) powered by stellar-mass black holes (BHs) with anisotropic X-ray emission (e.g. King et al. 2001) or with super-Eddington accretion rate/luminosity (e.g. Begelman 2002; King \\&{} Pounds 2003; Socrates \\&{} Davis 2006; Poutanen et al. 2007) or with a combination of the two mechanisms (e.g. King 2008). ULXs could also be associated with HMXBs powered by intermediate-mass BHs (IMBHs), i.e. BHs with mass $100\\,{}{\\rm M}_\\odot{}\\le{}m_{\\rm BH}\\le{}10^5\\,{}{\\rm M}_\\odot{}$ (see van der Marel 2004 for a review). IMBHs with mass larger than $100\\,{}{\\rm M}_\\odot{}$ may be required to explain the brightest ULXs (i.e. the $\\lesssim{}4$ ULXs with $L_{\\rm X}\\gtrsim{}10^{41}$ erg s$^{-1}$), those ULXs showing quasi-periodic oscillations (M82 X-1, see Strohmayer \\&{} Mushotzky 2003, and NGC 5408 X-1, see Strohmayer et al. 2007) and some of those that are surrounded by isotropically ionized nebulae (Pakull \\&{} Mirioni 2002; Kaaret, Ward \\&{} Zezas 2004). However, IMBHs are not needed to explain the observational properties of most of the ULXs (e.g. Gon\\c{c}alves \\&{} Soria 2006; Stobbart, Roberts \\&{} Wilms 2006; Copperwheat et al. 2007; Roberts 2007;  Zampieri \\&{} Roberts 2009). Thus, the objects classified as ULXs might actually be an inhomogeneous sample, including sources of different nature.  Most of ULXs are located in galaxies with a high star formation rate (SFR, e.g.  Irwin, Bregman \\&{} Athey  2004), although a small fraction (10--20 per cent, especially in the low-luminosity tail) might be associated with an old population (e.g. Colbert et al. 2004; Brassington, Read \\&{} Ponman 2005). The ULXs match the correlation between X-ray luminosity and SFR reported by various studies (Grimm, Gilfanov \\&{} Sunyaev 2003; Ranalli, Comastri \\&{} Setti 2003; Gilfanov, Grimm \\&{} Sunyaev 2004a,b,c; Kaaret \\&{} Alonso-Herrero 2008). Furthermore, the same studies indicate that the luminosity function of ULXs is the direct extension of that of HMXBs. Recent papers suggest a correlation between ULXs and low-metallicity environments, and propose that this may be connected with the influence of metallicity on the evolution of massive stars (Pakull \\&{} Mirioni 2002; Zampieri et al. 2004; Soria et al. 2005; Swartz, Soria \\&{} Tennant 2008). This scenario has been explored in detail by Mapelli, Colpi \\&{} Zampieri (2009, hereafter Paper~I) and by Zampieri \\&{} Roberts (2009), highlighting that a large fraction of ULXs may actually host massive ($\\sim{}30-80\\,{}{\\rm M}_{\\odot{}}$) stellar BHs formed in a low-metallicity environment. In fact, low-metallicity massive ($\\gtrsim{}40\\,{}{\\rm M}_\\odot{}$) stars lose only a small fraction of their mass due to stellar winds (Maeder 1992, hereafter M92; Heger \\&{} Woosley 2002, hereafter HW02; Heger et al. 2003, hereafter H03; Belczynski et al. 2010, hereafter B10) and can directly collapse (Fryer 1999; B10) into massive BHs ($25\\,{}{\\rm M}_\\odot{}\\le{}m_{\\rm BH}\\le{}80\\,{}{\\rm M}_\\odot{}$). Such massive BHs can power  most of the known ULXs without requiring super-Eddington accretion or anisotropic emission. Furthermore, their formation mechanism can explain the correlation between ULXs and SFR, and the fact that ULXs are preferentially found in low-metallicity  regions. In this Paper, we extend to a larger sample of galaxies the analysis reported in Paper~I, and we study the formation of massive BHs from the direct collapse of massive metal-poor stars and their possible connection with ULXs. In particular, we show that there is a correlation between the number of massive BHs formed in this scenario and the number of observed ULXs per galaxy. ",
        "conclusions": "Low-metallicity ($Z\\lesssim{}0.4\\,{}Z_\\odot{}$) massive ($\\gtrsim{}40\\,{}{\\rm M}_\\odot{}$) stars are expected to produce massive remnants ($25\\le{}m_{\\rm BH}/{\\rm M}_\\odot{}\\le{}80$, H03, B10) at the end of their evolution. Such massive BHs might power a large fraction of the observed ULXs in low-metallicity galaxies. In this Paper, we derived the number of massive BHs (N$_{\\rm BH}$) that are expected to form in a galaxy, via this mechanism, in the same time period that they could have a massive ($\\gtrsim{}15$ M$_\\odot{}$) donor companion. We find that N$_{\\rm BH}$ correlates  well with the observed number of ULXs per galaxy (N$_{\\rm ULX}$). The slope of such correlation does not depend significantly either on the IMF or on the adopted stellar evolution model. The IMF and the stellar-evolution models affect only the normalization of  N$_{\\rm BH}$ (the spread is generally less than a factor of 2). We stress that the stellar-evolution models adopted in this Paper neglect some important effects, such as the rotation and the possible influence of binary evolution. The final mass of the remnant is likely affected by the fact that the massive progenitor of the BH is fast rotating or that it resides in a binary, where additional mass loss is possible (see e.g. H03). Accounting for the probability that the progenitor of the massive BH is in a binary system likely introduces additional uncertainties.   In addition, the model considered here does not include the possibility of pair instability supernovae (PISNs). PISNs are predicted to occur at very low metallicity, in the case of very massive stars ($\\ge{}140$~M$_\\odot${}), and  lead to the complete disruption of the star (HW02; H03). Pulsational pair instability may occur also for smaller stellar masses ($100-140$~M$_\\odot{}$) and/or for larger metallicity, but it does not lead to a PISN, and a massive BH can form via direct collapse. Thus, PISNs probably do not play a role for stars with metallicity $Z\\gtrsim{}0.01\\,{}Z_\\odot$ and mass $<140$~M$_\\odot{}$ (H03). On the other hand, even assuming (as a strong upper limit) that all stars with mass $\\ge{}100\\,{}M_\\odot{}$ do not leave any remnant, due to a PISN, our estimates of N$_{\\rm BH}$ change by less than 10 per cent.   The model described in this Paper is consistent with the observed correlation between the number of ULXs and the SFR, as well as with the fact that ULXs are found preferentially in low-metallicity environments. Furthermore, this model is a natural extension of the scenario described by Grimm et al. (2003). In fact, Grimm et al. (2003) find a correlation between the SFR and the number of X-ray sources (not necessarily ULXs) per galaxy, and they explain it with the correlation between the SFR and the number of HMXBs powered by stellar BHs. In this Paper we suggest that this correlation still holds for the ULXs, as the ULXs (or most of them) can be powered by massive BHs formed by the collapse of  massive metal-poor  stars.   Our model predicts the existence of a dependence of N$_{\\rm ULX}$ on $Z$ (and the data  suggest it, too), when the dominant effect due to the SFR is removed. Unfortunately, the statistical uncertainty of such dependence  is still quite high, due to the dearth and to the inhomogeneity of the data.   In particular, there is a large inhomogeneity in the metallicity measurements. Moreover, the metallicity needed in our model is that of the molecular clouds before the pollution from the first supernovae associated with the parent clusters, as very massive stars collapse into BHs before such supernovae. Thus, the metallicity measured today is likely higher than the value we should consider in our model. Only for some types of galaxies, such as the ring galaxies, where the star-formation history has a clear connection with the geometry of the system, is it possible to measure a pre-starburst value of $Z$, suitable for our purposes. A possible way to reduce this problem and to check our model is to take new measurements of the local metallicity in the neighbourhoods of the observed ULXs, or even in the nebula associated with the ULXs (Ripamonti et al., in preparation)."
    },
    "1005/1005.4704_arXiv.txt": {
        "abstract": "We study the spectrum of high frequency radiation emerging from mildly dissipative photospheres of long-duration gamma-ray burst outflows. Building on the results of recent numerical investigations, we assume that electrons are heated impulsively to mildly relativistic energies by either shocks or magnetic dissipation at Thomson optical depths of several and subsequently cool by inverse Compton, scattering off the thermal photons of the photosphere. We show that even in the absence of magnetic field and non-thermal leptons, inverse Compton scattering produces power-law tails that extend from the peak of the thermal radiation, at several hundred keV, to several tens of MeV, and possibly up to GeV energies. The slope of the high-frequency power-law is predicted to vary substantially during a single burst, and the model can easily account for the diversity of high-frequency spectra observed by BATSE. Our model works in baryonic as well as in magnetically dominated outflows, as long as the magnetic field component is not overwhelmingly dominant. ",
        "introduction": "Understanding the origin of the prompt emission of Gamma-Ray Bursts (GRBs) has been hampered by two major challenges that have proven formidable obstacles despite more than four decades of studies. Radiating photons from an ultra-relativistic jet requires first a mechanism to convert at least some of the bulk kinetic energy into thermal motion of electrons, second a mechanism to radiate the electrons' energy into photons. These two processes are known as the problem of the dissipation mechanism and of the radiation process.  The standard model that has been developed over the years assumes that internal shocks are responsible for the dissipation (Rees \\& M{\\'e}sz{\\'a}ros 1994), while electrons gyrating in a shock-generated magnetic field produce the radiation via the synchrotron process (M{\\'e}sz{\\'a}ros et al. 1994; Piran 1999; Lloyd \\& Petrosian 2000). Both mechanisms are fraught with numerous problems.  Internal shocks assume that the outflow is released by the central engine with fluctuations in the Lorentz factor, so that different parts of the flow collide with each other, producing shocks that dissipate energy. Unfortunately, the internal shock mechanism has very little predictive power, since any behavior of the light curve can be explained by a suitable choice of the ejection history of the central engine. Since the physics of the emission from the engine is largely unconstrained, it is very hard to disprove internal shocks based on any observation. The only robust prediction of internal shocks as a dissipation mechanism is the low efficiency of the process (Kobayashi et al. 1997; Lazzati et al. 1999; Spada et al. 2000), due to the fact that internal shocks can dissipate only the energy associated with differential motions and not the energy associated to the bulk motion, which is much larger. Observations, instead, show that at least for a fraction of bursts, the efficiency of the prompt phase is 50 per cent, if not higher (Zhang et al.  2007). An alternative model that is becoming increasingly popular is magnetic dissipation in a Poynting dominated outflow (e.g., Thompson 1994; Spruit et al. 2001; Giannios \\& Spruit 2005). Even though magnetic dissipation could in principle provide very high efficiencies, it is, again, a very uncertain process and as such provides very little predictive power to allow for meaningful comparisons with observations.  Synchrotron radiation, with possible self-Compton components, has long been considered the best candidate to explain the prompt emission of GRBs. While synchrotron radiation can easily explain the non-thermal nature of the observed spectrum, more thorough scrutiny reveals several important problems. First, optically thin synchrotron emission has a very well defined limiting slope $\\alpha\\ge-1/3$ (where $F(\\nu)\\propto\\nu^{-\\alpha}$) due to the synchrotron radiation spectrum of a single electron (the so-called line of death of synchrotron emission). However, at least several cases of GRBs with $\\alpha<-1/3$ have been detected (Crider et al. 1997; Preece et al. 1998; Ghirlanda et al. 2002, 2003). Second, the typical slope of the low-frequency part of the GRB spectrum is $\\alpha=0$ (Kaneko et al. 2006), which is not a natural slope of synchrotron radiation from shock-accelerated electrons. Third, the high frequency spectrum has a highly variable slope $\\beta$, not easily explained by synchrotron radiation from shock accelerated electrons that should produce a fairly standard spectrum with $\\alpha\\sim1$ (as in the X-ray afterglows). Finally, for the prompt emission to be efficient, electrons are expected to cool fast (Ghisellini et al. 2000), but the typical slope of cooling electrons ($\\alpha=1/2$) is not observed in GRB spectra (Kaneko et al. 2006).  These well known difficulties of synchrotron emission have favored the study of various alternatives, including quasi-thermal Comptonization (Ghisellini \\& Celotti 1999), jitter radiation (Medvedev \\& Loeb 1999; Medvedev 2000; Workman et al. 2007; Morsony et al. 2008), bulk Compton (Lazzati et al. 2000); and photospheric emission (M{\\'e}sz{\\'a}ros \\& Rees 2000; M{\\'e}sz{\\'a}ros et al. 2002; Rees \\& M{\\'e}sz{\\'a}ros 2005; Pe'er et al. 2005, 2006; Giannios 2006; Giannios \\& Spruit 2007; Pe'er et al. 2007; Thompson et al. 2007; Ryde \\& Pe'er 2009, Lazzati et al. 2009). In this paper we expand the analysis of photospheric emission as a GRB prompt radiation mechanism, studying the conditions under which the photosphere can produce a high-frequency component with a non-thermal power-law shape up to tens of MeV. Photospheric radiation has two great advantages with respect to synchrotron emission: it does not require a dissipation mechanism and it naturally provides high efficiency. Photospheric radiation does not require a dissipation mechanism if the radiation is released before the full acceleration of the fireball and therefore at a stage when the fireball still contains a large fraction of internal energy. Numerical simulations have shown (Lazzati et al. 2009, hereafter LMB09) that this is indeed the case for long duration GRBs with a massive star progenitor. The photospheric efficiency of a typical GRB was found to be in good agreement with the observations (LMB09; Zhang et al. 2007). The obvious weakness of photospheric radiation is that it is customarily assumed to be thermal, therefore lacking the prominent non-thermal tails observed in GRBs.  The possibility of adding non-thermal tails to the photospheric spectrum has been investigated before. Pe'er et al. (2006) showed that continuous dissipation and/or internal shocks at moderate optical depths give rise to a Comptonized spectrum with a flat energy spectrum $\\nu F(\\nu)\\propto\\nu^0$. Giannios (2006, see also Giannios \\& Spruit 2007) performed a similar study and reached analogous conclusions, focusing on magnetic reconnection as the dissipation mechanism. Here we argue that a dissipation mechanism that manages to be so continuous as to keep the electron temperature at a constant equilibrium value under the intense cooling of IC scattering is very unlikely. A more realistic assumption is that sub-photospheric electron heating takes place as a series of one or more episodes of impulsive acceleration and subsequent cooling. We show that in this case a non-thermal high-frequency tail is produced, characterized by a slope that is sensitive to a combination of various parameters and which is therefore expected to be highly variable during the prompt emission of GRBs.  This paper is organized as follows: in \\S~2 we discuss the origin of the non-thermal tails, in \\S~3 we discuss the propagation of the non-thermal spectrum in the optically thick sub-photospheric plasma, and in \\S~4 we present Monte Carlo calculations of the model. We summarize and discuss our results in \\S~5. ",
        "conclusions": "We have presented a method by which a primary blackbody spectrum of a GRB outflow photosphere can be Comptonized into a non-thermal high-frequency power-law tail. In this scenario, the bulk of the GRB prompt emission is thermal photospheric and the non-thermal tail contains a relatively small fraction of the radiation energy. The majority of the energy of the radiation is therefore released before it is converted into bulk kinetic energy, solving the problem of identifying a dissipation mechanism with high efficiency (LMB09).  The non-thermal tail is produced by thermal electrons accelerated to mildly relativistic energies by an intermittent dissipation mechanism. The electrons are allowed to cool through IC interactions with the photon field between acceleration events. This dissipation mechanism is not required to have an efficiency larger than a few per cent, since the power-law tail only contains a fraction of the total energy of the radiation.The spectral slope of the high energy radiation depends on the typical Lorentz factor of the accelerated electrons, on the peak frequency of the thermal photon spectrum, and on the ratio of the photon to lepton density. Differently from previous work (Pe'er et al. 2005, 2006; Giannios 2006) we assume that the electrons are accelerated intermittently and not continuously and that the electrons are allowed to cool among acceleration events. In addition, we assume that the energy density in radiation is larger than that in the relativistic electrons, a situation that is easily realized in jets born inside massive progenitor stars (LMB09).  Compared to the standard optically thin synchrotron, this mechanism can explain the steep low-frequency slopes observed in the early phases of some bright bursts (Crider et al. 1997; Preece et al. 1998; Ghirlanda et al. 2002, 2003) and the transient thermal bumps detected is several events (Ryde 2005; Ryde \\& Pe'er 2009; Ryde et al. 2010). Bursts for which no dissipation takes place should not display any non-thermal features. A complete discussion of the low-frequency non-thermal tail is however beyond the scope of this paper and will be presented elsewhere. Explaining a typical low-frequency tail $F(\\nu)\\propto\\nu^0$ is indeed a challenge for any model and not only for the photospheric scenarios presented here and elsewhere (M{\\'e}sz{\\'a}ros \\& Rees 2000; M{\\'e}sz{\\'a}ros et al. 2002; Rees\\& M{\\'e}sz{\\'a}ros 2005; Pe'er et al. 2005, 2006; Giannios \\& Spruit 2007; Pe'er et al. 2007; Thompson et al. 2007; Ryde \\& Pe'er 2009, LMB09).  Radiation from dissipative photospheres has been investigated before (Pe'er et al. 2005, 2006; Giannios 2006; Giannios \\& Spruit 2007; Beloborodov 2010). Giannios (2006) and Giannios \\& Spriut (2007) concentrate on the case of a continuous energy injection or ``slow heating'', different from our impulsive acceleration assumption. Pe'er et al. (2005, 2006) consider instead two possible scenarios: ``slow heating'' and impulsive heating by internal shocks. They find that the bulk of the photons are shifted in energy once a steady state electron population is attained. They calculate that, for a fairly wide parameter range, the comoving peak frequency is of few tens keV. Attaining such a steady state population requires however a large optical depth and as a consequence we do not find the same results. To understand the reason for this, consider an injection episode at an optical depth of a few. Due to IC scattering, the hot electrons are efficiently cooled and their energy given to a small fraction of photons that are shifted to high frequencies. For small optical depth, the energized photons reach the photosphere with only negligible losses and decouple from the electron population. The energy given to the electrons by the dissipation event is therefore given to a small number of photons that produce a power-law tail as shown in the figures. In order to obtain a steady state electron population able to exchange energy effectively with the radiation field, a large number of scattering per photon is required, so that the high-frequency photons can return their energy to the electron population. Large optical depths are therefore required for electron and photon populations to settle in a steady state configuration dictated by the balance of heating and cooling. Such optical depths are much larger than the ones envisioned in this work. In any case, a fraction of the dissipated energy is left in the electron population that thermalizes at a temperature higher than the one before the dissipation. For the parameters adopted in this paper the effect of the increased electron temperature is negligible because the increase of the electrons temperature is very small for the optical depths considered.  Finally, Beloborodov (2010) investigated the radiation produced by a well-defined dissipation mechanism: nuclear and Coulomb collisions within a baryonic outflow with a substantial population of neutrons (see also Beloborodov 2003; Rossi et al. 2006). He finds that a typical spectrum with a high-frequency slope $F(\\nu)\\propto\\nu^{-1.4}$ is obtained. Differently to our result, the non-thermal spectrum is due to the presence of non-thermal leptons.  A characteristic of any photospheric model is that the radiation is released at a small radius $r_{\\rm{ph}}\\sim10^{10-13}$~cm (e.g. Rees \\& M{\\'e}sz{\\'a}ros 2005). The main consequence of such a small radius is that the compactness of the region is large (Pe'er \\& Waxman 2004) and therefore photon-photon interactions result in the production of electron-positron pairs. This can have two consequences. If the newly created pairs do not outnumber the original electrons, the only consequence is the presence of a sharp cutoff in the spectrum at $h\\nu=511$~keV in the comoving frame ($\\sim100$~MeV observed). If, on the other hand, the pairs outnumber the original electrons, a new photosphere is created, typically at a radius 3 to 10 times larger than the original photosphere (Rees \\& M{\\'e}sz{\\'a}ros 2005). The scenario we developed in this paper is not dependent on whether the photosphere is due to the original electrons or to pairs, the only difference being that a pair photosphere has a lower photon to lepton ratio and produces therefore a steeper high-frequency spectrum. In addition, since the number of pairs and the number of photons are not independent, we showed that a pair photosphere with at least several energization events would produce spectra with a typical slope $1<\\beta_{\\rm{sat}}<1.5$, in good agreement with observations (e.g., Kaneko et al. 2006).  The Fermi satellite has recently shown that at least some bright GRBs have spectra extending well into the GeV regime (Abdo et al. 2009abc; 2010; Ghisellini \\& Ghirlanda 2010; Granot et al. 2010). For any reasonable combination of parameters, GeV radiation cannot be produced at the photosphere, since it would be immediately absorbed by photon-photon collisions to produce pairs. In this scenario, the GeV emission can be either internal, i.e. produced by dissipation above the photosphere (see \\S~2.2; Toma et al. 2010) or external, due to the interaction of the outflow with the interstellar medium (Kumar \\& Barniol Duran 2010; Ghirlanda et al. 2010; Ghisellini et al. 2010).  In terms of predictive power, the most notable prediction of this model is the presence of substantial variability of the high-frequency spectrum due to the dependence of the spectrum on parameters that are expected to be highly variable such as, e.g., the peak frequency of the thermal spectrum in the comoving frame (LMB09). Such variability might be however hard to detect since it is expected to have short time-scales and would therefore be averaged out in a time-integrated spectrum. The dependence of the spectral slope on the peak frequency of the spectrum opens also the possibility of correlations between the peak frequency and the spectral slope at high frequencies. Since, however, the slope depends on two more parameters (the ratio of photons to electrons and the typical Lorentz factors of the electrons) it is not obvious that such correlations should be visible in experimental data."
    },
    "1005/1005.3915_arXiv.txt": {
        "abstract": "{We obtained \\emph{Herschel} PACS and SPIRE images of the thermal emission of the debris disk around the A5V star $\\beta$\\,Pic.  The disk is well resolved in the PACS filters at 70, 100, and 160\\,$\\mu$m. The surface brightness profiles between 70 and 160\\,$\\mu$m show no significant asymmetries along the disk, and are compatible with 90\\% of the emission between 70 and 160\\,$\\mu$m originating in a region closer than 200\\,AU to the star.  Although only marginally resolving the debris disk, the maps obtained in the SPIRE 250 -- 500\\,$\\mu$m filters provide full-disk photometry, completing the SED over a few octaves in wavelength that had been previously inaccessible. The small far-infrared spectral index ($\\beta = 0.34$) indicates that the grain size distribution in the inner disk ($<$200\\,AU) is inconsistent with a local collisional equilibrium. The size distribution is either modified by non-equilibrium effects, or exhibits a wavy pattern, caused by an under-abundance of impactors which have been removed by radiation pressure. } ",
        "introduction": "The $\\beta$\\,Pic disk, discovered by IRAS \\citep{aumann:1984}, was the first debris disk to be directly imaged in scattered light \\citep{smith:1984}.  It is seen close to edge-on and extends in the optical out to 95\\arcsec, corresponding to 1800\\,AU \\citep{larwood:2001}.  $\\beta$\\,Pic (A5V) is one of the closest \\citep[19.44\\,$\\pm$\\,0.05\\,pc,][]{vanleeuwen:2007} and youngest debris disks.  The estimated age \\citep[12 Myr,][]{zuckerman:2001} significantly  exceeds typical timescales for the survival of pristine circumstellar dust grains \\citep[e.g.,][]{fedele:2010}, hence continuous replenishment of the dust, presumably through collisions of planetesimals, is needed. The closeness of the object ensures that it can also be spatially resolved at long wavelengths: \\citet{holland:1998} resolved the disk at 850\\,$\\mu$m and \\citet{liseau:2003} at 1200\\,$\\mu$m.  Optical and near-infrared observations of the inner part ($<$100AU) of the disk yield evidence of asymmetries such as warps and density contrasts, which may relate to the presence of planetesimals \\citep{kalas:1995,pantin:1997,mouillet:1997,heap:2000, telesco:2005}. \\citet{lagrange:2009B} imaged a possible companion at a projected distance of 8\\,AU from the star.  Images of $\\beta$\\,Pic in scattered stellar light directly detect small grains and indirectly larger grains that produce the smaller ones through collisions.   The grain-size distribution can be quantitatively constrained from the spectral energy distribution (SED) of the disk, in the infrared and (sub)mm domains. The spectral index of the SED at the longest wavelengths \\citep{liseau:2003,nilsson:2009} is inferred to be fairly low, which according to modeling \\citep{draine:2006,natta:2007,ricci:2010} can be interpreted as a deficit of small grains.  In this paper, we present far-infrared imaging of the $\\beta$\\,Pic debris disk in six \\emph{Herschel} photometric bands between 70 and 500\\,$\\mu$m.  These bands cover for the first time the long-wavelength side of the peak in the thermal emission of the disk, and the large aperture of the telescope enables us to resolve the disk at far-IR wavelengths for the first time.  With these data, we measure the surface brightness profiles of the disk and readdress the issue of the grain-size distribution in the inner 200\\,AU. ",
        "conclusions": "We have presented images of the $\\beta$\\,Pic debris disk in six photometric bands between 70 and 500\\,$\\mu$m using the PACS and SPIRE instruments. We resolve the disk at 70, 100, 160, and 250\\,$\\mu$m. The images at 70--160\\,$\\mu$m show no evidence of asymmetries in the far-infrared surface brightness along the disk of $\\beta$\\,Pic. The observed profiles are compatible with 90\\% of the emission originating in a region within a radius of 200\\,AU from the star. The disk-integrated photometry in the six \\emph{Herschel} filters provides a far infrared SED with small spectral index $\\beta\\,\\approx\\,0.34$, which is indicative of a grain size distribution that is inconsistent with a local collisional equilibrium. The size distribution is modified by either non-equilibrium effects, or exhibits a wavy pattern, caused by the under-abundance of impactors that are small enough to be removed by radiation pressure."
    },
    "1005/1005.2408_arXiv.txt": {
        "abstract": "We propose a novel method to estimate \\mm, the ratio of stellar mass (\\mstar) to black hole mass (\\mbh) at various redshifts using two recent observational results: the correlation between the bolometric luminosity of active galactic nuclei (AGN) and the star formation rate (SFR) in their host galaxies, and the correlation between SFR and \\mstar\\ in star-forming (SF) galaxies. Our analysis is based on \\mbh\\ and \\lbol\\ measurements in two large samples of type-I AGN at z$\\simeq$1 and z$\\simeq$2, and the measurements of \\mm\\ in 0.05$<$z$<$0.2 red galaxies. We find that \\mm\\ depends on \\mbh\\ at all redshifts. At z$\\simeq$2, \\mm$\\sim$280 and $\\sim$40 for \\mbh=$10^8$ and \\mbh=$10^9$\\msun, respectively. \\mm\\ grows by a factor of $\\sim$4-8 from z$\\simeq$2 to z=0 with extreme cases that are as large as 10--20. The evolution is steeper than reported in other studies, probably because we treat only AGN in SF hosts. We caution that estimates of \\mm\\ evolution which ignore the dependence of this ratio on \\mbh\\ can lead to erroneous conclusions. ",
        "introduction": "\\label{sec:intro}  Study of the co-evolution of Active Galactic Nuclei (AGN) and their host galaxies provides important clues about the growth of super massive black holes (SMBHs) and the star formation (SF) history of the Universe. In the local universe one finds a tight correlation between the SMBH mass (\\mbh) and the mass of the bulge of its host, \\mbul\\ (Marconi \\& Hunt, 2003; H\\\"{a}ring \\& Rix 2004, hereafter HR04), or alternatively with the stellar velocity dispersion \\sigs\\ (Ferrarese \\& Merrit 2000; Gebhardt \\et 2000). Typically \\mbul/\\mbh$\\simeq$500--1000. $\\mbul/\\mbh$ must have been smaller at high redshift. For example, $\\mbh\\sim10^{9-10}\\,\\Msun$ are often observed at z=3--6 (e.g. Netzer 2003; Shemmer \\et 2004; Fan \\et 2006), yet galaxies that are 500-1000 times more massive are never observed at z$>$0.5 and are very rare even at z$\\simeq$0.  The evolution of the $\\mbh-\\sigs$ relationship has been studied in numerous papers. Examples are Shields \\et (2003) who find no evolution up to z$\\simeq$2, and Woo \\et (2008) who suggest that $\\mbul/\\mbh$ has increased by a factor of $\\sim$3 since z=0.6. The uncertainties in all such measurements are very large due to the difficulties in measuring \\sigs\\ or \\mbul\\ in high redshift AGN hosts. While measuring \\mbul\\ in high-redshift galaxies is difficult, the total stellar mass, \\mstar\\ is easier to obtain. This is achieved by multi-band spectral energy distribution (SED) fitting, which is used to constrain \\ml. However, obtaining \\mm\\ for AGN hosts is severely limited by the problematic subtraction of the bright point-like continuum in type-I AGN and the estimation of \\mbh\\ in type-II AGN.  Several recent studies used deep imaging to estimate \\mm\\ (or $L_{\\rm Host}/M_{\\rm BH}$) in z$\\sim$0.5--3 AGN by resolving the host galaxy emission and by careful PSF modelling (e.g. Kukula \\et 2001; Peng \\et 2006; Kotilainen \\et 2007; Bennert \\et 2010). The measured host luminosity was translated to \\mstar\\ by \\textit{assuming} a certain \\ml\\ ratio. A common assumption is that AGN hosts are ``red and dead'', with a stellar population which evolves passively from a high formation redshift, e.g. $z_{\\rm form}\\simeq$5. A detailed study of this type, including a summary of many earlier findings, is given in Decarli et al. (2010; hereafter D10) who find that \\mm\\ evolves following $\\mmwp\\propto\\,z^{-0.28}$. The assumption of non-SF AGN hosts may describe some objects, but many studies find that hosts of luminous AGN often contain much younger stellar populations (e.g., Kauffmann \\et 2003; Jahnke \\et 2004; Silverman \\et 2009; Merloni \\et 2010, and references therein).    In this \\textit{Letter} we suggest a novel method to estimate the evolution of \\mm. Our approach makes use of the known relationships between star formation rate (SFR) and \\mstar\\ in SF galaxies (SFGs), and between the bolometric luminosity of AGN (\\lbol) and the SFR in their hosts. We describe these relationships in \\S\\ref{sec:sfr_corr} and use them to estimate \\mm\\ in SF galaxies at z$\\simeq$0.15, \\zandz\\ in \\S\\ref{sec:results}. In \\S\\ref{sec:discussion} we discuss the implications to the co-evolution of SMBHs and their hosts. Throughout this work we assume a standard $\\Lambda${\\sc CDM} cosmology with $\\Omega_{\\Lambda}$=0.7, $\\Omega_{M}$=0.3 and $H_{0}$=70\\,\\kms\\,Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:discussion}  The analysis presented here suggests strong evolution in \\mm\\ up to z=2, much steeper than what is suggested in other studies. In particular, our mean values of \\mm\\ at \\zandz\\ are smaller than the ones reported in D10, which is the most up-to-date compilation of such works. There are three main reasons for these differences: \\begin{figure} \\centering \\includegraphics[width=6.2cm]{Fig3_TN10.eps} \\caption{The distributions of \\mm\\ at different redshift, for the three \\mbh\\ sub-groups discussed in the text.} \\label{fig:MM_hist_3z_3m} \\end{figure} \\begin{enumerate} \\item As already mentioned, most earlier works estimate \\mstar\\ by assuming a passive evolution from a high formation redshift. This is obviously an over-simplification of the evolution of most galaxies and significant epochs of SF at z$<$4 (e.g. D07; E07; Drory \\& Alvarez 2008; van Dokkum \\et 2010, and references therein). Accounting for younger stellar populations would result in lower \\mstar\\ and thus lower \\mm\\ (see D10 and Peng \\et 2006). \\item The D10 sample represents the \\textit{minority} of the AGN population, as hinted by two biases. First, the majority of the D10 sources lie in the top 15\\% of the \\mbh\\ distributions corresponding to their redshift. Second, the selection criteria for the D10 HST observations could have been biased towards large, resolved galaxies with large \\mstar. \\item The mean \\mbh\\ in high redshift AGN samples is systematically larger than the corresponding low redshift \\mbh. For example, in our large sample of red galaxies, 99\\% of the sources show \\mbh$<10^{8.8}$\\msun\\ while 20\\% of the z$\\simeq$1 and 57\\% of the z$\\simeq$2 AGN have larger \\mbh. All the z$\\simeq$2 sources in D10 have \\mbh$\\gtsim10^{8.7}$\\msun. Such objects should only be compared with local galaxies which host BHs that are at least as massive. As Fig.~\\ref{fig:MM_vs_Mbh_all} shows, this corresponds to \\mm=100--200, instead of the commonly used $\\sim$700. \\end{enumerate} In conclusion, while our work applies to most AGN, the D10 sample probably represents the remaining sources.  The results presented here point to a scenario where many galaxies have to increase their mass by factors of 4-8 (2-4) since z$\\simeq$2 (z$\\simeq$1). The growth factors for the most massive BHs are not well determined since the number of very massive galaxies in the local universe is not large enough to reliably extend the results of Fig.~\\ref{fig:MM_vs_Mbh_all} beyond $\\mbh\\simeq10^9\\msun$. These numbers represent the requirement for the galaxies to \\textit{over}-grow their SMBHs by the above factors. This seems to be consistent with models which suggest that the high mass SMBHs observed at z$\\simeq2$ could have accumulated most of their mass by that redshift (Marconi \\et 2004). On the other hand, it may be in contradiction with at least some scenarios linking AGN activity to the shut-down of SF in their host galaxies (see Somerville \\et 2008; Cattaneo \\et 2009 and references therein).   While a full discussion of the various growth scenarios of \\mstar\\ is beyond the scope of this \\textit{Letter}, we comment briefly on some of these ideas. Major galaxy mergers would increase both \\mstar\\ and \\mbh, either through starbursts and gas accretion in ``wet mergers'' or the possible coalescence of the two SMBHs involved in ``dry mergers''. However, theoretical and observational arguments (e.g. Lotz \\et 2008; Genel \\et 2009) suggest a low rate of such events for z$<2$ galaxies. Thus, major mergers cannot change \\mm\\ by more than a factor of $\\sim2-3$ between z=2 and z=0. Small  ``dry mergers'' may help. For example, Naab, Johansson \\& Ostriker (2009) show that present-day massive red ellipticals gain $\\sim40\\%$ of their mass through accretion of smaller companions since z$\\sim2$. Intense SF in outer parts of galaxies due to external source of cold gas which does not find its way to the centre (e.g. van Dokkum \\et 2010), is another possibility. More possibilities and more references are discussed in Benson \\& Devereux (2009). The \\mm\\ distributions presented here suggest an increase in \\mm\\ by factors beyond what is suggested in many theoretical studies.  Finally, we comment on the possibility that the suggested evolution of \\mm\\ could be due to two wrong assumptions. First, many more AGN hosts may not lie on the SF sequence or may not obey the \\lbol-SFR correlation used here. This is unlikely to be the case at low redshift, where SDSS type-II AGNs are used. However, the selection of at least some of the most luminous, high redshift high-\\lbol\\ sources in Fig.~\\ref{fig:Lbol_Lsf} may be biased towards high FIR luminosity, high SFR AGN hosts (e.g. Zheng \\et 2009) in particular if these are found in mergers that are not part of the SF sequence. \\textit{Herschel} observations of well-defined AGN samples are likely to resolve this issue. Second, the Drory \\& Alvarez (2008) work shows a decline in SFR at the high-\\mstar\\ end for z$<$2 galaxies. Thus, some of our \\lbol-based SFR estimates might be associated with considerably larger values of \\mstar. The E07 and D07 SF sequences do not show such a decline.  We conclude that there is a steep evolution in \\mm\\ from z$\\simeq2$ to z=0 for SF AGN hosts. This trend is barely consistent with some, but not all galaxy and BH evolution models. We have also demonstrated the crucial importance of considering different \\mbh\\ groups separately when evaluating the \\mm\\ evolution.  \\vspace*{0.4cm}  \\noindent We thank an anonymous referee for very useful comments on the manuscript. We thank Ido Finkelmann for fruitful discussions; Vincenzo Mainieri for allowing us to use zCOSMOS results ahead of publication; Roberto Decarli and Emanuele Daddi for sending us their data; Niv Drory and Samir Salim for their thoughtful comments; and the MPA/JHU team for making their SDSS catalogues available to the public. This study makes use of data from the SDSS (http://www.sdss.org/collaboration/credits.html). Funding for this work has been provided by the Israel Science Foundation grant 364/07.  \\vspace*{ -0.6cm}"
    },
    "1005/1005.0916_arXiv.txt": {
        "abstract": "SDSS\\,J094857.3+002225 is a very radio-loud narrow-line Seyfert 1 (NLS1) galaxy. Here, we report our discovery of the intranight optical variability (INOV) of this galaxy through the optical monitoring in the \\emph{B} and \\emph{R} bands that covered seven nights in 2009. Violent rapid variability in the optical bands was identified in this RL-NLS1 for the first time, and the amplitudes of the INOV reaches 0.5 mag in both the \\emph{B} and \\emph{R} bands on the timescale of several hours. The detection of the INOV provides a piece of strong evidence supporting the fact that the object carries a relativistic jet with a small viewing angle, which confirms the conclusion drawn from the previous multi-wavelength studies. ",
        "introduction": "Narrow-line Seyfert 1 galaxies (NLS1s) are a special and interesting group of active galactic nuclei (AGNs). They show narrow optical Balmer emission lines [FWHM(H$\\beta$) $<$ 2000 km s$^{-1}$], weak [\\ion{O}{3}]$\\lambda$5007 emission ([\\ion{O}{3}]/H$\\beta<3$), strong \\ion{Fe}{2} emission, and soft X-ray excess \\citep{pog00,ost85,sul00,bol96,bol97}.  NLS1s show remarkable radio-loud/radio-quiet bimodality \\citep{lao00}. Only $7\\%$ of NLS1s are radio-loud objects \\citep{zho03,kom06}. The fraction is much smaller than that found in QSOs. Very radio-loud NLS1s (RL-NLS1s,  $R>100$) are even much fewer ($\\sim2.5\\%$) \\citep{kom08}, where the radio loudness $R$ is commonly defined as the flux ratio of radio to optical at $\\lambda4400$ \\citep{kel89}. So far, it is still a puzzle why RL-NLS1s are so scarce. At present, the origin of RL-NLS1s is also still poorly understood.    A few efforts have been made in the past few years to understand the nature of RL-NLS1s. Yuan et al. (2008) found that the broadband spectra of some RL-NLS1s are similar to those of high-energy-peaked BL Lac objects, and suggested that some of them may be BL Lac objects actually. Basing upon the recent observation taken by \\emph{Fermi} satellite, some RL-NLS1s display a hard X-ray component suggesting the presence of relativistic jets on the line of sight \\citep{fos09,abd09a}.  The presence of the relativistic jets motivates us to search for intranight optical variability in some RL-NLS1s, because of the well-known beaming effect (e.g., Wagner \\& Witzel 1995). \\citet{kom06} argued that SDSS\\,J094857.3+002225 is a right candidate for searching for RL-NLS1s with beaming effect. The object is a very radio-loud NLS1 at $z=0.585\\pm0.001$. The reported radio loudness derived from the radio flux at 5 GHz ranges from 194 to 1982 \\citep{zho03,kom06,abd09a}. It is in the CRATES catalog as a flat-spectral radio source \\citep{hea07}. The simultaneous observations taken by both \\emph{Swift} and \\emph{Fermi} also suggest that the broadband spectral energy distribution is similar to those of flat-spectral radio quasars \\citep{abd09a,zho03}. Recent photometry from the Guide Star Catalogs 2.21 is $B_{\\mathrm{J}}(\\mathrm{GSC2.21})=18.83$ mag \\citep{zho03}.    Previous studies revealed multi-wavelength variabilities in the object at timescales from day to year. Previous radio observations indicate its fluctuation in the radio band on the timescale from weeks to years \\citep{abd09a}.  \\citet{zho03} also said that the object shows long-term variability in both the radio and optical bands. The amplitude of the variation in the radio can be $60\\%$ within a year. The long-term variability amplitude may be about 1 mag in the optical band. The latest multi-wavelength campaign carried out by \\citet{abd09b} discovered an optical variability on day timescales. Dramatic flux variabilities in both X-rays and radio 37 GHz were also found in the study.  In this Letter, we report an optical monitor for the RL-NLS1 SDSS\\,J094857.3+002225. The monitor was designed to search for intranight optical variability (INOV) in the object. The INOV should be detected if the object indeed hosts a relativistic jet beaming toward the observers. ",
        "conclusions": "The particular RL-NLS1 galaxy SDSS\\,J094857.3+002225 was monitored in optical bands by NAOC 80 cm TNT telescope to search for its INOV phenomenon. Our optical monitoring indeed provides clear evidence for the presence of INOV in both the \\emph{B} and \\emph{R} bands in the object. The object exhibits optical variability  not only on the timescale of a week, but also on several hours. The detection of the INOV indicates that the object contains a relativistic jet on the line of sight of an observer, which confirms the conclusion drawn from the high-energy observations (e.g., Abdo et al. 2009a, 2009b) and from the inverted radio spectrum and high brightness temperature (Zhou et al. 2003).   SDSS\\,J094857.3+002225 is particular for its observational properties. On the one hand, its optical spectrum with strong \\ion{Fe}{2} emission is typical of NLS1s. The narrow H$\\beta$ emission yields a relatively low black hole (BH) mass $\\sim 4.0\\times10^7\\ M_\\odot$ and a high Eddington ratio (Zhou et al. 2003). On the other hand, some observational behaviors are characteristic of blazars with the relativistic jets close to the line of sight, such as the INOV detected here, flat radio spectrum, high brightness temperature, and variable $\\gamma$-ray emission (see citations in Section 1).  So far, outstanding RL-NLS1s with blazer-like radio emission have been revealed by multi-wavelength observations in several cases including the object SDSS\\,J094857.3+002225. We refer the reads to Yuan et al. (2008) for a brief summarization. With the successful launch of \\emph{Fermi} satellite, $\\gamma$-ray emission was detected in four RL-NLS1s, including the object studied here, which suggests the presence of fully developed jets in these objects \\citep{abd09c}. The authors argued that the four RL-NLS1s may form a new class of $\\gamma$-ray AGNs because of their small BH masses, large Eddington ratios, and possibly disk-like morphology of the host galaxies.   The intrinsic mechanism of RL-NLS1s is an attractive field. There are two possible models for interpreting RL-NLS1s. The first one is the inclination model \\citep{ost85,kom06,wha06,kom08}. The model suspects that the (at least a fraction of) RL-NLS1s are preferentially viewed pole-on. The observed narrow width of the Balmer emission lines could be resulted from small inclination if the broad-line region (BLR) is constrained to a plane \\citep{wil86,bia04}. In fact, there is some evidence supporting a flat BLR in some RL-AGNs \\citep{jar06,sul03}. In this scenario, the BH mass is largely underestimated in these objects since the current available estimation of the BH mass of AGN from single-epoch spectroscopic observation comes from an assumption of an isotropic distribution of the broad-line clouds with random orbital inclinations \\citep{kas05,pet04,ben07}. Although the inclination model sounds reasonable because it is able to shift the location of RL-NLS1s on the $R$--$M_{\\mathrm{BH}}$ plane to the massive BH end \\citep{lao00,lac01}, the massive BHs are not supported by the lack of massive bulges in several cases in which the host galaxies can be resolved.   The second is the accretion mode model \\citep{kom06,wha06,kom08}. The RL-NLS1s with small BH masses are accreting close to or even above the Eddington limit. Low-mass BHs may lead to narrow emission lines when Keplerian velocities are considered mostly \\citep{wha06}. The accretion is thought related to the radio emission. Different accretion modes may result in different phenomena and maybe can explain the differences in the radio loudness. Accretion processes are known related to the spin of the accreting BHs. So the rapid spin of BHs may also affect the radio loudness of NLS1s \\citep{kom06}. It is possible that accretion mode combining with BH spin can explain the nature of the RL-NLS1s.   Although our monitor indicates that the extremely high radio emission in RL-NLS1 SDSS\\,J0948\\\\57.3+002225 is mainly contributed from the beamed non-thermal jet with a small viewing angle (can also be found in the aforementioned other studies), more information is needed in the future to investigate the origin of the relativistic jet."
    },
    "1005/1005.0127_arXiv.txt": {
        "abstract": " ",
        "introduction": "As an alternative to inflation, the curvaton model \\cite{Lyth:2001nq} is a well-motivated proposal for explaining the observed scale-invariant primordial density perturbation in the framework of inflation (and some pioneering works on this mechanism have been considered in Refs. \\cite{Mollerach:1989hu, Linde:1996gt, Enqvist:2001zp}). This model is based on the inflationary scenario of multiple fields, in which two fields are required at least. In this model the universe is composed of radiation decayed from the inflaton and the curvaton field. During inflation, the curvaton is subdominant and provides entropy (isocurvature) perturbation, and afterwards, the entropy perturbation can be converted to curvature perturbation as long as the curvaton decays into radiation before primordial nucleosynthesis. After the curvaton decays, the universe enters the standard thermal history, and then the primordial curvature perturbation leads to the formation of the large scale structure of our universe\\cite{Lyth:2002my, Lyth:2003ip}.  Recently, a curvaton scenario realized in the frame of stringy inflationary models was presented in Ref. \\cite{Li:2008fma}. The idea of brane inflation proposed in \\cite{Dvali:1998pa}, is recently successfully realized by virtue of warped compactifications\\cite{Kachru:2003si}. In the stringy landscape a number of probe branes are allowed to move in warped throats, and under the description of supergravity (SUGRA) their dynamics is described by a Dirac-Born-Infeld (DBI) action in the relativistic limit. When applied to cosmology, it was found that a relativistically moving D3-brane in a deformed AdS throat is able to drive inflation without slow-roll at early universe\\cite{Silverstein:2003hf, Chen:2004gc}. This idea was later extended into the scenario of multiple brane inflation appeared in Refs. \\cite{Cai:2008if, Cai:2009hw} (see \\cite{Cline:2005ty, Piao:2002vf} for earlier studies under slow-roll approximation). A distinguish feature of the model in this type is that a large positive non-local non-Gaussianity can be obtained due to an enhancement of a small sound speed for the perturbation\\cite{Alishahiha:2004eh, Chen:2006nt}, which is in contrast to the prediction of a canonical single-field inflation model\\cite{Maldacena:2002vr}. Established on the scenario of multiple brane inflation\\cite{Cai:2008if, Cai:2009hw} and motivated by the model of {\\it spinflation} \\cite{Easson:2007dh}, an explicit construction of the curvaton brane model was recently realized in Ref. \\cite{Zhang:2009gw} in which angular degrees of freedom had been introduced to perform a relativistic rotation in a warped throat.   The investigation of non-Gaussianities is crucial in distinguishing the curvaton scenario from usual inflation models. A key prediction of a general curvaton model is that a sizable local non-Gaussianity can be obtained of which the value is mainly decided by the occupation of curvaton energy $\\Omega_{\\chi}$ after inflation characterized by a transfer efficiency parameter $$r=\\frac{3\\Omega_{\\chi}}{4-\\Omega_{\\chi}}$$ which is usually calculated on the hypersurface of the curvaton decay. As a consequence, the decay mechanism of a curvaton field plays an important role in determining the detailed information of primordial perturbation seeded by the curvaton field \\cite{Sasaki:2006kq}. It was pointed out in Ref. \\cite{Sasaki:2006kq} that at third order, the primordial curvature perturbation under the assumption of sudden curvaton decay is nicely consistent with that obtained in the case of a non-instantaneous decay in the limits of small $r$ and $r\\rightarrow1$; however, their bispectra and trispectra deviate from the others in the middle band of the allowed value regime of $r$. The implications of the non-Gaussian features of curvature perturbation on the curvaton scenario was studied in \\cite{Huang:2008ze, Huang:2008bg, Gong:2009dh}.  In the present paper we extend the investigation of Ref. \\cite{Sasaki:2006kq} relating non-Gaussian features of primordial curvature perturbation to the curvaton decay mechanism into a more general curvaton scenario. Specifically, we consider a model of curvaton brane. We study the generation of curvature perturbation and non-Gaussianity originating from the vacuum fluctuations on the curvaton brane in cases of a sudden decay on the slice of uniform curvaton density, and then compare their observational signatures with the standard curvaton model to illustrate the significance of the decay process of the curvaton field. Moreover, since the model of curvaton brane involves a sound speed parameter which describes the propagation of the curvaton fluctuations during inflation, it has interesting implications to cosmological observations as well. Note that in our scenario, the inflaton field can be realized by any string realization of inflation which is independent of the curvaton field. Thanks to the virtue of curvaton scenario, the detailed realization of inflation does not affect the calculation for curvaton perturbation, as long as the inflaton perturbation itself is negligible.  Our paper is organized as follows. In Section 2, we point out that the non-Gaussian features in a curvaton mechanism crucially depend on the curvaton decay mechanism. We suggest the curvaton could decay on a slice of uniform curvaton density, and non-Gaussian fluctuation evaluated on this slice is mainly determined by a combination of the equation-of-state (EoS) of the curvaton and a generalized transfer efficiency parameter which will be shown in the main text. In Section 3, we study the model of curvaton brane as a specific example. We firstly describe the background evolution of a curvaton brane model in relativistic limit throughout the inflationary period and the process of curvaton decay. Afterwards, we develop an analytic investigation of the curvature perturbation and the non-Gaussianity generated in this model by solving the perturbation equation of the curvaton brane order by order. In Section 4 we consider a class of plausible curvaton brane dominant eras and derive the nonlinearity parameters characterizing non-Gaussian distribution of the primordial curvature perturbation in these cases respectively. Finally, we present concluding remarks in Section 5. ",
        "conclusions": "In this paper, we have calculated the nonlinear primordial curvature perturbation in the curvaton scenario using a generalized $\\delta{N}$ formalism. If curvaton decay does not occur on the slice of uniform total energy density but on the slice of uniform curvaton density instead, we find that the dynamics of the nonlinear fluctuations behaves different from the usual scenario. Specifically, we consider a model of curvaton brane which can provide a theoretical realization of the curvaton decay we expected. In the frame of this model, we have presented a full analysis on the sizes and shapes of its bispectra and trispectra respectively. We arrive at an important conclusion that the generation of nonlinearities is sensitive to the mechanism of curvaton decay, which is mainly determined by the EoS $w$, the density occupation $\\Omega_{\\chi}$ of the field $\\chi$ at curvaton decay and its sound speed $c_s$ during inflation. Explicitly speaking, our results show that the nonlinearity parameter characterizing the second order perturbation $f_{NL}$ is proportional to $1/\\tilde{r}$ in which the transfer efficiency relies on the EoS $w$ and the density occupation $\\Omega_{\\chi}$ of the curvaton. For the equilateral shape, the size of $f_{NL}$ is amplified by a large value of the factor $1/c_s^2$, and therefore when the curvaton brane moves relativistically the shape of non-Gaussianity is mainly of equilateral type; however, at the squeeze limit the nonlinearity parameter $f_{NL}$ decouples from the sound speed and its size is qualitatively consistent with the result obtained in a usual canonical curvaton model. Moreover, for the third order fluctuations, we find the nonlinearity parameter $g_{NL}$ is enhanced by the factor $1/c_s^4$ in equilateral shape and decouples in local limit, and in the simplest example its size is proportional to $1/\\tilde{r}^2$ which is quite different from the usual scenario.  As an end, we would like to highlight the importance of our study in this paper. The mechanism of curvaton decay after inflation could determine the relation of the primordial curvature perturbation and the decay hypersurface. This process is rather robust and should be considered in any specific curvaton model, and the example of curvaton brane analyzed in the current work is a good illustration to emphasize its importance. Additionally, the signatures of curvaton decay imprinted on the primordial curvature perturbation could provide a new window to explore the combination of early universe physics and astronomical observations.   \\medskip"
    },
    "1005/1005.5228_arXiv.txt": {
        "abstract": "We investigate damping and growth times of the quadrupolar $f$-mode for rapidly rotating stars and a variety of different polytropic equations of state in the Cowling approximation. This is the first study of the damping/growth time of this type of oscillations for fast rotating neutron stars in a relativistic treatment where the spacetime degrees of freedom of the perturbations are neglected. We use these frequencies and damping/growth times to create robust empirical formulae which can be used for gravitational wave asteroseismology. The estimation of the damping/growth time is based on the quadrupole formula and our results agree very well with Newtonian ones in the appropriate limit. ",
        "introduction": "\\label{sec:introduction}  During the birth of a proto-neutron star or the merging of two older compact stars, violent non-radial oscillations may be excited, resulting in the emission of significant amounts of gravitational radiation \\cite{Andersson:2011lr}. The detection of gravitational waves from oscillating neutron stars will allow the study of their interior, in the same way as helioseismology provides information about the interior of the Sun. It is expected that the identification of specific pulsation frequencies in the observational data will reveal the true properties of matter at densities that cannot be probed today by any other experiment. In this paper, we present new empirical relationships for mode-frequencies and damping times of the quadrupolar $f$-mode for rapidly rotating neutron stars, extending previous studies which deal with the non-rotating case \\cite{1996PhRvL..77.4134A,1998MNRAS.299.1059A,2001MNRAS.320..307K}.   These original suggestions about gravitational wave asteroseismology have been supported by many complementary works which studied specific features of oscillation spectra for various compact objects, such as typical neutron stars \\cite{1999MNRAS.310..797B,2002PhRvD..65b4010S,2004PhRvD..70l4015B,2005PhRvL..95o1101T,2006PhRvD..74l4025T,2007PhRvD..76d3003C,2007MNRAS.381..151W,Lau:2010lr}, but also for strange \\cite{2003PhRvD..68b4019S,2004PhRvD..69h4008S,Benhar:2007lr} or superfluid stars \\cite{2001PhRvL..87x1101A,2002PhRvD..66j4002A}. More recently, it has also been suggested that one may use asteroseismology to find the imprints of scalar or even vector components of gravity \\cite{2004PhRvD..70h4026S,2005PhRvD..71l4038S,2009PhRvD..79f4033S,2009PhRvD..80f4035S}. It should be noted here that all previous studies  \\cite{Andersson:2011lr,1996PhRvL..77.4134A,1998MNRAS.299.1059A,2001MNRAS.320..307K,1999MNRAS.310..797B,2002PhRvD..65b4010S,2004PhRvD..70l4015B,2005PhRvL..95o1101T,2006PhRvD..74l4025T,2007PhRvD..76d3003C,2007MNRAS.381..151W,Lau:2010lr,2003PhRvD..68b4019S,2004PhRvD..69h4008S,Benhar:2007lr,2001PhRvL..87x1101A,2002PhRvD..66j4002A,2004PhRvD..70h4026S,2005PhRvD..71l4038S,2009PhRvD..79f4033S,2009PhRvD..80f4035S} have been performed for non-rotating relativistic stars. The treatment of rotation was always a problem in general relativity and thus the majority of the studies for the oscillation spectra of fast-rotating compact stars was done mainly in Newtonian theory which gives only qualitative answers.  Since stellar oscillations may become unstable in the presence of rotation, there was  an increased interest during the last decade or so to study the dynamics of rotating stars, also thanks to the discovery of the $r$- and $w$-mode instability \\cite{1998ApJ...502..708A,1998ApJ...502..714F,1998PhRvL..80.4843L,1999ApJ...510..846A,Kokkotas:2004lr}. Still, the majority of these studies have been performed in Newtonian theory \\cite{1999PhRvD..59d4009L,1999ApJ...521..764L,1999PhRvD..60f4006L} while there are only a few works in which GR has been used; mainly in the so called {\\em slow-rotation approximation}. The slow-rotation approximation was successfully applied to study various aspects of the $r$-mode instability \\cite{2001IJMPD..10..381A,2001MNRAS.328..678R}, effects of uniform and differential rotation on the oscillation spectrum \\cite{Ferrari:2004qy,2007PhRvD..75f4019S,2008PhRvD..77b4029P} and on the crustal modes \\cite{2008MNRAS.384.1711V}.  As it has been originally suggested  by Chandrasekhar \\cite{1970PhRvL..24..611C} and verified by Friedman \\& Schutz \\cite{1978ApJ...221..937F,1978ApJ...222..281F}  certain non-axisymmetric pulsation modes may grow exponentially in rotating stars; this is due to the emission of gravitational waves and is called  CFS instability. Exploring this type of instability in rapidly rotating stars turned out to be very difficult. In linear perturbation theory for example, rapid rotation was never treated properly until recently; almost all formulations of the relevant perturbation equations were prone to numerical instabilities either at the surface or along the rotation axis of the neutron star. Thus, it was not surpising that the first results for the oscillations of rapidly rotating stars were derived  using evolutions of the non-linear equations \\cite{2004MNRAS.352.1089S,2006MNRAS.368.1609D,2006PhRvD..74f4024K,2008PhRvD..77f4019K}. Still all these studies were purely axisymmetric and thus the effects of rotation on the spectra was present only for very high rotation rates. Rotational instabilities are driven by non-axisymmetric modes and thus these first 2D calculations where not of much use for their study.  In the last two years there was significant progress in the study of non-axisymmetric perturbations of rapidly rotating neutron stars. For the first time it was possible to calculate in GR the oscillation spectra of fast rotating relativistic stars by using the linearized form of the fluid equations.  Thus the effect of fast rotation on  $f$- and $r$-modes has been demonstrated  while the  critical points for the onset of the $f$-mode (CFS) instability have been derived \\cite{2008PhRvD..78f4063G}. In addition it has been demonstrated that there is a way to derive empirical relations connecting the oscillation frequencies with the rotation of the stars. This study has been recently extended  to $g$-modes \\cite{2009PhRvD..80f4026G} and even more recently has been expanded to study the oscillation spectra of fast and differentially rotating neutron stars \\cite{2010PhRvD..81h4019K}.  It should be noted that the previous results have been derived using the so-called Cowling approximation where the spacetime is assumed to be frozen. This approximation is very good for the estimation of the spectra of $r$- and $g$-modes but it gives only qualitatively good results for the $f$-mode. Moreover, using non-linear codes it became possible for the first time to study the complete problem \\cite{PhysRevD.81.084055}, i.~e.~the non-axisymmetric stellar oscillations of fast rotating stars without the constraints of the Cowling approximation. The results are in qualitative agreement with those found in \\cite{2008PhRvD..78f4063G} and for the critical point for the onset of the $f$-mode instability with the studies presented in \\cite{1998ApJ...492..301S}.  The next step for gravitational wave asteroseismology is to use additional information about the damping times to construct model-independent relations which allow for a robust determination of stellar key parameters. The damping time $\\tau$ of the potentially CFS-unstable branch in the high rotation regime for example can be approximated very accurately by using a simple relation of the form $\\tau_0/\\tau = \\sgn{(\\sigma_i)}\\,0.256(\\sigma_i/\\sigma_0)^4$, where $\\tau_0$ and $\\sigma_0$ are the damping time and mode frequency of the nonrotating model respectively, $\\sgn{(x)}$ is the signum function and $\\sigma_i$ is the actual mode frequency in the inertial frame.  The structure of the paper is as follows. In Section \\ref{sec:problem_setup} we give an essential overview about our method of computing mode-frequencies and damping times of the $f$-mode. We then show the results of our simulations in Section \\ref{sec:empirical_relations}, where we present empirical relations which can potentially be used to estimate masses, radii and angular frequencies of rapidly rotating neutron stars. A more elaborate discussion about the numerical procedure, the equations of state and background models used in this study as well as a consistency check can be found in the Appendix. ",
        "conclusions": "\\label{sec:summary} In this work we demonstrated how one can do gravitational wave asteroseismology by using the frequencies and possibly the damping/growth times of the emitted waves from oscillating and rapidly rotating relativistic stars.  This is possible by the empirical relations that we have derived and which connect the frequencies and the damping/growth times of the oscillation modes with the stellar characteristics, i.e. with the mass, radius and rotation rate. We have actually shown that for polytropic equations of state of varying stiffness one can create very robust formulae connecting the observable frequency and damping times with the quantities like rotation frequency, average density and/or compactness.  We have shown on a few examples how one can use the empirical formulae in order to derive the stellar parameters. In a realistic situation when an $f$-mode will be excited, it will be possible to detect the signal at least from galactic sources if the mode is CFS-stable and at least from sources in the Virgo cluster if it is unstable \\cite{Lai:1995fk,Ou:2004qy,Shibata:2004uq,2010KWK}. This will be possible with the sensitivity of the advanced Virgo and LIGO detectors \\cite{Acernese:2006fk,Abbott:2009qy} and probably even more feasible with the next generation gravitational wave telescopes such as ET (Einstein Telescope) \\cite{2010CQGra..27h4007P,Andersson:2011lr}.  The ``weak'' point of the whole procedure relies in the approximate calculation of both frequencies and damping times. As we already mentioned we have neglected the spacetime perturbations and thus there is a systematic quantitative but not qualitative error in all data. Thus it is expected that the coefficients in relations \\eqref{eq:sigmaStableBranch}, \\eqref {eq:sigmaUnstableBranch} , \\eqref{eq:fitForDampTimeUnstable} and \\eqref{eq:fitForDampTimeStable} will be affected by a proper treatment of the spacetime degrees of freedom. However, it is believed that these changes will not alter the results significantly since the relations for frequencies and damping times are normalized with their corresponding values in the non-rotating limit. These values will absorb most of the differences when compared with the correct results in the presence of spacetime perturbations but of course, this is an issue that has to be addressed properly in future work. An additional outcome of this analysis will be the frequencies and damping times of the $w$-modes.  Finally, the empirical relations found in this study have been derived for polytropes of varying stiffness which are able to mimic the global properties of realistic equations of state; see for example EoS A and EoS II which are polytropic fits to tabulated EoS. Realistic hot equations of state are the best candidates for newly born neutron stars and have a higher chance of becoming unstable prior mutual friction completely suppresses any instability \\cite{Andersson:2009lr} but currently our code is unable to perform time-evolutions of rapidly rotating neutron stars for generic tabulated data. Preliminary studies regarding this issue are promising but in a very early stage."
    },
    "1005/1005.2064_arXiv.txt": {
        "abstract": "We estimated absolute shifts of Fe~I and Fe~II lines from Fourier-transform spectra observed in solar active regions. Weak Fe~I lines and all Fe~II lines tend to be red-shifted as compared to their positions in quiet areas, while strong Fe~I lines, whose cores are formed above the level $\\log \\tau_5\\approx-3$   (about 425~km), are relatively blue-shifted, the shift growing with decreasing lower excitation potential. We interpret the results through two-dimensional MHD models, which adequately reproduce red shifts of the lines formed deep in the photosphere. Blue shifts of the lines formed in higher layer do not gain substance from the models. ",
        "introduction": "We continue the study initiated by us in the papers \\cite{4,5} concerned with the effect of small-scale magnetic fields on solar granulation. The study is based on the observations made by P. Brandt with a Fourier-transform spectrometer (FTS) at the McMath telescope (Kitt Peak National Observatory, USA). The advantages of FTS observations are their high quality and a possibility to observe almost concurrently a large number of lines. Areas close to the solar disk center were observed.  Here we investigate in great depth absolute shifts of Fe~I and Fe~II lines in low-activity areas (plages) with different integral magnetic fluxes. This is the first time that such an investigation covers a large number of lines formed at various levels in the solar photosphere so that effects of small-scale magnetic fields on height distribution of photospheric velocity field may be studied in detail.  The present state of the investigations on small-scale solar magnetic fields is elucidated in review  \\cite{20} by Solanki. Determinations of absolute shifts of absorption lines in active areas are not numerous, they are overviewed in  \\cite{1}.  Paper  \\cite{7} by Brandt and Solanki, similar to  \\cite{4,5} and the present paper, is also based on the FTS observations, but only 32 lines were used there to study line parameter variations and 19 lines to determine absolute shifts. Those 19 lines were strong Fe~I lines only. The conclusion was made that the lines in active areas had red shifts with respect to their positions in the quiet photosphere. The shifts depend on magnetic field strength (filling factor) and are, on the average, 0.22 pm (about 120~m/s) for lower bisector sections. This result agrees qualitatively with the findings of the studies by Cavallini et al.  \\cite{8}, Livingston  \\cite{17}, Immerschitt and Schroter \\cite{15}.  At the same time Cavallini et al.  \\cite{9} dealing with absolute shifts of four lines formed at different photospheric levels obtained a quite different result --- shifts of line cores with respect to their positions in the quiet photosphere depend on the region of line core formation: the weak Fe~I line at $\\lambda$  614.92~nm displays a red shift, while the strong Ca~I line at  $\\lambda$ 616.22~nm is blue-shifted. A similar result was obtained by Keil et al.  \\cite{16} for the strong Fe~I $\\lambda$ 543.45~nm line with a small lower excitation potential (1.01~eV), the core of this line is formed in the upper photospheric layers. Absolute shifts of the line were found to depend on magnetic field polarity, but when the shifts are averaged without regard for local field signs, the line core has a blue shift of about 150~m/s.  Evidence of different behavior of velocity fields with height in the photosphere in active and quiet areas is found also in the analysis of spectral observations with high spatial resolution made by Hanslmeier et al.  \\cite{13,14}. If this is the case, lines formed at different levels in the atmosphere should display different absolute shifts.  Thus, the problem needs to be thoroughly investigated on the basis of an extended sample of spectral lines. This is the prime objective of the present study. ",
        "conclusions": "FTS observations were used to estimate possible absolute shifts of Fe~I and Fe~II lines in active regions. We conclude that weak Fe~I lines and all Fe~II lines are likely to be red-shifted with respect to their positions in the quiet photosphere, while strong Fe~I lines with low $EPL$s, their cores being formed above the level $\\log \\tau_5\\approx-3$ ($\\approx$425~km), have more violet shifts.  The results were interpreted on the basis of 2-D MHD models that reproduce quite adequately the red shifts of the lines formed deep in the photosphere, but violet shifts of the lines formed in the upper photospheric layers are not confirmed within the scope of these models.  \\vspace{1.cm} {\\bf Acknowledgements.} We wish to thank the National Optical Astronomical Observatory/National Solar Observatory (Tucson, Arizona, USA) and especially J. Brault and B.~Graves for support in the FTS observations; S. Solanki for filling factor calculations; R.~Rutten for assistance in the organization of this investigation. The study was partially financed by the Joint Foundation of the Government of Ukraine and International Science Foundation (Grant No. K11100)."
    },
    "1005/1005.3987_arXiv.txt": {
        "abstract": "White dwarfs inspiraling into black holes of mass $\\MBH\\simgt 10^5M_\\odot$ are detectable sources of gravitational waves in the LISA band.  In many of these events, the white dwarf begins to lose mass during the main observational phase of the inspiral.  The mass loss starts gently and can last for thousands of orbits.  The white dwarf matter overflows the Roche lobe through the $L_1$ point at each pericenter passage and the mass loss repeats periodically.  The process occurs very close to the black hole and the released gas can accrete, creating a bright source of radiation with luminosity close to the Eddington limit, $L\\sim 10^{43}$~erg~s$^{-1}$. This class of inspirals offers a promising scenario for dual detections of gravitational waves and electromagnetic radiation. ",
        "introduction": "One of the goals of the Laser Interferometer Space Antenna (LISA) mission is to detect gravitational waves from compact stellar objects spiraling into massive black holes, a class of events called extreme-mass-ratio inspirals (EMRIs) (e.g., Hils \\& Bender 1995; Gair et al. 2004; Barack \\& Cutler 2004; see Hughes 2009 and Sathyaprakash \\& Schutz 2009 for recent reviews).  Of particular interest are sources of coincident gravitational and electromagnetic radiation. Besides providing unique information on the nature of the event, such dual detections will lead to a new version of Hubble diagram that is based on the gravitational distance measurements (e.g. Bloom et al. 2009; Phinney 2009).  Inspirals into black holes of masses $\\MBH\\sim (10^5-10^6)M_\\odot$ produce gravitational waves in the frequency band where LISA is most sensitive.  Normal stars are tidally disrupted well before they approach such black holes, and therefore discounted as possible LISA sources (however, see Freitag 2003).  Inspirals of compact objects are guaranteed sources of gravitational waves, however most of them are not promising for dual detections.  In particular, stellar-mass black hole or neutron star inspirals are not expected to generate bright electromagnetic signals.  Only inspiraling white dwarfs (WDs) offer a possibility for dual detection (Menou et. al. 2008; Sesana et al. 2008). WDs can be tidally disrupted very close to the black hole and then could create a transient accretion disc with Eddington luminosity.  Estimated EMRI rates are high enough for observations with LISA (e.g. Phinney 2009). The expected fraction of WD inpirals among all EMRIs depends on the degree of mass segregation in galactic nuclei, which favors stellar-mass black holes over white dwarfs in the central cluster.  The abundance of white dwarfs also depends on the details of stellar evolution, which are not completely understood. A fraction of WD inpirals as large as $\\sim 10\\%$ has been suggested (e.g. Hopman \\& Alexander 2006a).  The orbital parameters of WD inspirals are also uncertain.  EMRIs form when a compact object is captured onto a tight orbit whose evolution is controlled by gravitational radiation rather than random interactions with other stars in the central cluster around the massive black hole. Two main channels exist for EMRI formation: capture of single stars and capture of binary systems (e.g., Hils \\& Bender 1995; Sigurdsson \\& Rees 1997; Ivanov 2002; Gair et al. 2004; Hopman \\& Alexander 2005, 2006a,b, 2007; Miller et al. 2005; Hopman 2009). In the single-capture scenario, the shrinking orbit can retain a significant eccentricity until the end of inspiral.  In contrast, the binary-capture scenario leads to nearly circular orbits (Miller et al. 2005).  In this paper, we focus on WD inspirals that are not completed because the star is tidally disrupted before its orbit becomes unstable.  We argue that the WD begins to lose mass very gently and, for thousands of orbital periods, this process resembles accretion through the $L_1$ point in a binary system rather than a catastrophic disruption.  This offers a possibility of {\\it simultaneous} observation of the inspiral by LISA and traditional, optical and X-ray telescopes.  Previous work on tidal deformation of a WD orbiting a massive black hole focused on two extreme regimes: (i) weak deformation was studied analytically using perturbation theory (e.g., Rathore et al. 2005; Ivanov \\& Papaloizou 2007), and (ii) strong deformation leading to immediate disruption was simulated numerically (e.g. Kobayashi et al. 2004; Rosswog, Ramirez-Ruiz \\& Hix 2009).  The regime considered in the present paper is different from both cases explored previously.  It involves an extended phase of strong deformation with small mass loss, which we call `tidal stripping' below.  The mass of the WD remains almost unchanged during this phase and continues to emit gravitational waves. Even a small orbital eccentricity, e.g. $e=0.01$, implies that tidal stripping occurs only near the pericenter of the orbit, during a small fraction of each orbital period.  Thus, the mass loss is expected to be {\\it periodic}, possibly leading to a periodic electromagnetic signal from the inspiral. ",
        "conclusions": "The standard picture of EMRI envisions an orbit that gradually shrinks due to gravitational radiation until its pericenter $r_p$ reaches $\\rmin$ where the orbit becomes unstable and plunges into the black hole. The main observational phase of the inspiral is when $r_p$ decreases from $\\sim 2\\rmin$ to $\\rmin$. For WDs with mass $M\\sim M_\\odot$ inspiraling into a black hole with $\\MBH\\sim 10^5M_\\odot$ the main observational phase lasts $\\sim 10^5$ orbital periods, which may be comparable to one year, depending on the orbital eccentricity. We argued in Section~2 that the inspiraling WDs can experience an extended period of slow mass loss during observations by LISA.  To summarize, the tidal stripping begins very gently because the pericenter of the orbit drifts inward slowly, by a tiny fraction $\\alpha\\sim 10^{-5}$ in one orbital period. This leads to many repeated episodes of small mass loss. The process occurs at radius $r_0$ (eq.~\\ref{eq:r0}), comparable to $\\rmin$.  The star is strongly deformed by the tidal forces near $r_0$ but it barely touches its Roche lobe for a small fraction of the orbital period. As a result, the star loses a small amount of mass through the $L_1$ point at each pericenter passage. To our knowledge, this regime of tidal stripping was not explored by direct hydrodynamical simulations.  Our crude estimates suggest two phases of stripping: the first phase lasts $N_1\\sim 10^3$ orbits until the star loses $\\sim 1$ per cent of its mass.  Then the mass loss accelerates: the decrease in $M$ (and the corresponding increase in the WD radius) implies that the star overfills the Roche lobe more and more with every pericenter passage. As a result, the star loses the remaining 99 per cent of its mass in a few hundreds of additional orbits. The mass lost in each individual episode during these last hundreds of orbits behaves as $\\dM\\propto (N_\\star-N)^{-\\xi}$ where $N_\\star-N\\ll N_\\star$ is the number of orbits remaining to complete disruption and $\\xi\\sim \\psi/(\\psi-1)$.  According to our estimates, tidal stripping operates for days or weeks, creating a relatively long-lived source of gas.  The gas can accrete onto the black hole and produce significant electromagnetic radiation together with the gravitational waves observed by LISA.  The radiation source may become bright before the mass loss spoils the standard EMRI template for the gravitational-wave signal.  The gas produced by tidal stripping moves on nearly Keplerian orbits that are initially close to the WD orbit. The gas probably leaves the star with relative velocity $\\sim v_{\\rm esc}=(2GM/R)^{1/2}$ and its orbital energy differs from that of the star by a small fraction $\\sim v_{\\rm esc}/v\\sim(Mr_0/\\MBH R)^{1/2} \\sim\\gamma^{-1/2}(M/\\MBH)^{1/3}\\sim 1/30$. After $\\sim 30$ orbital periods, the differential rotation of the gas has stretched it into a ring around the black hole.  A non-zero eccentricity of the donor orbit will create an eccentric ring.  It should viscously spread and accrete onto the black hole.  A small mass-loss fraction $x$ can be a huge source of gas for accretion.  If most of the stripped matter is accreted by the black hole, the accretion rate is $\\dot{M}\\sim \\dM/P$. It exceeds the Eddington value $\\dot{M}_{\\rm Edd}\\sim 10^{23}(\\MBH/10^5M_\\odot)$~g~s$^{-1}$ after $N\\sim 10$ orbits since the beginning of tidal stripping, well before the final disruption of the WD. The accretion timescale in the viscous ring can be estimated as $\\tacc\\sim \\alpha_v^{-1}(H/r)^{-2}P$ where $P$ is the WD orbital period, $\\alpha_v=0.01-0.1$ is the viscosity parameter and $H$ is the thickness of the ring (e.g. Shakura \\& Sunyaev 1973).  $H/r\\sim 1$ is expected for super-Eddington accretion, which leads to $\\tacc\\sim (10-100)P$.  This suggests that a bright source with Eddington luminosity $L_{\\rm Edd}\\sim 10^{43}$~erg~s$^{-1}$ is created quickly, in less than 1 day after the beginning of tidal stripping. For a typical distance to such LISA sources, $d\\sim 100$~Mpc, the accretion ring should be detectable with optical and X-ray telescopes provided its approximate location on the sky is known.  LISA is expected to localize EMRIs within $\\sim 10$~deg$^2$ (Barack \\& Cutler 2004). For massive black-hole mergers, the localization information will be available weeks to months prior to the final coalescence (Kocsis et al. 2007, 2008; Lang \\& Hughes 2008), and the localization expectations for WD EMRIs are similar (S. Drasco, J. Gair, I. Mandel, E. Porter, private communications), giving sufficient time for simultaneous optical and X-ray observations during inspiral.  Mass transfer in WD inspirals is special as it creates a long-lived source of gas very close to the black hole horizon, $r_0-r_g\\sim r_g$. As a result, the donor orbit is generally not confined to a plane, if the black hole rotation is significant.  Besides, the orbit will experience fast precession.  Therefore, the freshly stripped gas may collide with the previously released gas and generate shocks (e.g. Evans \\& Kochanek 1989).  The resulting pattern of accretion may be complicated and needs careful study.  An intriguing feature of tidal stripping is the periodic supply of gas. It may leave a fingerprint on the observed luminosity, modulating it with the WD orbital period $P$. The modulation of $\\dot{M}$ may create a detectable oscillation in the luminosity from the accreting ring, even though the accretion timescale $\\tacc\\sim (10-100)P\\gg P$, tends to reduce the amplitude of modulation. The shock emission from collisions between the periodic flow from the $L_1$ point and the gas accumulated around the black hole may be strongly modulated.  The description of mass transfer in this paper is greatly simplified. The possibility of many repeated mass-transfer episodes is robust, but the exact rate of tidal stripping and the dynamics of accretion need to be explored with dedicated numerical simulations.  The great potential that such events hold for joint detections of gravitational and electromagnetic radiation provides motivation for the effort. The joint detection would let us witness, in real-time and unprecedented detail, the slow tidal stripping of a WD followed by its complete disruption. Note that LISA observations are expected to provide the mass and spin of the black hole, as well as the details of the inspiral orbit. This can be used to model in detail the hydrodynamics of accretion.  The scenario discussed in this paper assumes $r_0>\\rmin$ and uses a semi-Newtonian description for the WD orbit. A fully relativistic model will be needed to accurately evaluate the parameter space for such events.  The relativistic effects are especially important if the black hole is rapidly rotating -- then $r_0$ and $\\rmin$ will depend on the black hole mass $\\MBH$, its spin parameter $a_s$, the eccentricity of the orbit, and the angle between the orbital angular momentum and the angular momentum of the black hole. The competition between tidal disruption and gravitational capture by rotating black holes was investigated for parabolic orbits in Beloborodov et al. (1992). In a broad range of $\\MBH$, the fate of a star approaching the black hole depends on the orbit orientation and can be either disruption or capture. For inspirals with $r_g<r_0<\\rmin$, tidal stripping does not occur. Instead, $r_p$ reaches $\\rmin$ and the orbit loses stability before any mass is lost by the WD. Then the star is crushed by tidal forces as it falls into the black hole. This strong and immediate disruption is different from the gentle stripping considered here, suggesting rich phenomenology of WD inspirals."
    },
    "1005/1005.3039_arXiv.txt": {
        "abstract": "Spacetime curvature plays the primary role in general relativity but Einstein later considered a theory where torsion was the central quantity.  Just as the Einstein-Hilbert action in the Ricci curvature scalar $R$ can be generalized to $f(R)$ gravity, we consider extensions of teleparallel, or torsion scalar $T$, gravity to $f(T)$ theories. The field equations are naturally second order, avoiding pathologies, and can give rise to cosmic acceleration with unique features. ",
        "introduction": "}  Acceleration of the cosmic expansion is one of the premier mysteries of physics and may be the clearest clue to properties of gravity beyond general relativity.  Extensions to gravity have been considered by making the action a function of the spacetime curvature scalar $R$ or other curvature invariants, by coupling this Ricci scalar to a scalar field, by introducing a vector field contribution, and by using properties of gravity in higher dimensional spacetimes.  Here we take a wholly different path, avoiding the curvature completely, although our results will end up with interesting relations to each of the above mentioned theories.  Rather use the curvature defined via the Levi-Civita connection, one could explore the opposite approach and use the Weitzenb{\\\"o}ck connection that has no curvature but instead torsion.  This has the interesting property that the torsion is formed wholly from products of first derivatives of the tetrad, with no second derivatives appearing in the torsion tensor.  In fact, this approach was taken by Einstein in 1928 \\cite{ein28,eineng}, under the name ``Fern-Parallelismus'' or ``distant parallelism'' or ``teleparallelism''.  It is closely related to standard general relativity, differing only in terms involving total derivatives in the action, i.e.\\ boundary terms.  In this paper, we investigate extensions where a scalar formed from contractions of the torsion tensor is promoted to a function of that scalar.  This parallels the concept of extension of the Ricci scalar $R$ in the Einstein-Hilbert action to a function $f(R)$, which has attracted much attention in recent years as a way to explain acceleration of the universe \\cite{durrer,sotiriou}.  The generalized $f(T)$ torsion theory has the advantage, however, of keeping its field equations second order due to the lack of second derivatives, unlike the fourth order equations (at least in the metric formulation) of $f(R)$ theory that can lead to pathologies. ",
        "conclusions": "\\label{sec:concl}  The class of $f(T)$ gravity theories is an intriguing generalization of Einstein's ``new general relativity'', taking a curvature-free approach and instead using a connection with torsion.  It is analogous to the $f(R)$ extension of the Einstein-Hilbert action of standard general relativity, but has the advantage of second order field equations.  We have also seen that it can be related to the form of modifications to the Friedmann equations due to higher dimensional gravity theories such as DGP.  It is also related to scalar-tensor gravity.  Writing the gravitational action as \\beq S=\\int d^4x\\,|e|\\,\\{T+f(T)+(T-A)\\,[1+f_A(A)]\\}\\,, \\eeq one can view the last term as a Lagrange multiplier term and find an equivalent scalar-tensor theory with $A=T$ and an effective potential \\beqa V_{\\rm eff}(\\psi)&=&\\frac{T}{1+f_T}-\\frac{T+f}{(1+f_T)^2} \\\\ \\psi&=&-\\ln(1+f_T)\\,. \\eeqa Furthermore, Einstein originally introduced teleparallelism to obtain a vector field component of the field equations \\cite{ein25,eineng}, intending to unify gravity and electromagnetism.  Recently, interest has grown in vector fields, ``Einstein aether theories'', as a way to obtain cosmic acceleration \\cite{vectorde}.  These theories can also give modifications to the field equations involving functions of $H^2$, i.e.\\ $T$ (see, e.g., \\cite{zuntz}).  Indeed they can be viewed as closely related to torsion theories (see \\cite{hayash1,hayash2} for details).  Thus, torsion theories can unify a number of interesting extensions of gravity beyond general relativity.  In investigating the nature of gravitation, we may find that Einstein presaged the acceleration of the universe not only through the cosmological constant but through a generalization of ``Einstein's other gravity''."
    },
    "1005/1005.3808_arXiv.txt": {
        "abstract": "The Planck experiment will soon provide a very accurate measurement of Cosmic Microwave Background anisotropies. This will let cosmologists determine most of the cosmological parameters with unprecedented accuracy. Future experiments will improve and complement the Planck data with better angular resolution and better polarization sensitivity. This unexplored region of the CMB power spectrum contains information on many parameters of interest, including neutrino mass, the number of relativistic particles at recombination, the primordial Helium abundance and the injection of additional ionizing photons by dark matter self-annihilation. We review the imprint of each parameter on the CMB and forecast the constraints achievable by future experiments by performing a Monte Carlo analysis on synthetic realizations of simulated data. We find that next generation satellite missions such as CMBPol could provide valuable constraints with a precision close to that expected in current and near future laboratory experiments. Finally, we discuss the implications of this intersection between cosmology and fundamental physics. ",
        "introduction": "Starting with COBE's groundbreaking detection of microwave background fluctuations \\cite{cobe}, the past two decades have seen dramatic improvements in measurements of the microwave background temperature fluctuations (see e.g. \\cite{altri} and \\cite{wmap0}). Planck's highly anticipated temperature power spectrum measurements (see \\cite{planck}) will further advance this program and produce significantly improved constraints on cosmological parameters.  While Planck's measurement of the anisotropy power spectrum to multipoles $\\ell \\sim 2000$ will extract most of the information in primordial temperature fluctuations, ongoing and planned ground-based and balloon-based experiments are exploring two important open frontiers: (a) the measurement of extremely ($\\le5'$) small-scale temperature and polarization fluctuations \\cite{small} and (b) the search for primordial B-modes, the distinctive signature of gravitational waves from inflation, on large scales \\cite{bmodesgw}.  For example, balloon-borne experiments such as EBEX \\cite{ebex} and SPIDER \\cite{spider} will improve the measurements of CMB polarization while ground based telescopes such as the Atacama Cosmology Telescope (ACT) \\cite{ACTPol} and the South Pole Telescope \\cite{spt} will extend temperature and polarization measurements to smaller, sub-arcminute, angular scales. Proposals for next generation CMB satellites such as CMBPol \\cite{CMBPol} or B-POL \\cite{bpol} are under evaluation from American and European space agencies.  What will we learn from measuring CMB temperature and polarization fluctuations on small-scales? The amplitude of temperature and polarization  fluctuations is determined by several different physical effects: (1) the amplitude of primordial fluctuations, (2) the evolution of the ionization fraction of the universe at $z > 1200$, which determines the sound speed for acoustic fluctuations, (3) the evolution of the ionization fraction of the universe at $z < 1200$, which determines the thickness of the surface of last scatter and (4) the transition from radiation to matter domination. Moreover, while small-scale CMB fluctuations are initially pure $E$ mode, gravitational lensing rotates $E$ modes into $B$ modes \\cite{cmblensing}. By measuring the pattern of small scale $E$ and $B$ modes, cosmologists will be able to determine the large-scale convergence field, a direct measure of the integrated density fluctuations between redshift $z=1100$ and $z=0$ (see e.g. \\cite{hirata},\\cite{okamotohu}). The convergence power spectrum is particularly sensitive to density fluctuations at $z \\sim 2$, an important complement to planned optical lensing measurements that probe the evolution of density fluctuations in the $z < 1$ universe.  The goal of this paper is to quantify the cosmological information that could come from these new datasets. This is important for several reasons. First, while there have been many studies of the future cosmological constraints from Planck, very few papers have investigated the constraining power of combinations of future CMB datasets from different sources. Second, as we will describe in the next sections, we will consider a large set of parameters focusing on those that mainly affect the \"damping tail\" of the CMB angular spectrum. We consider additional parameters such as the total neutrino mass $\\sum m_{\\nu}$ (which affects the growth of structure in the late universe), the number of additional relativistic neutrino species $N_{\\nu}^{eff}$ (which changes the matter-radiation epoch), and possible changes in the recombination process due to changes in the fractional helium abundance $Y_p$, dark matter self-annihilation processes, and variations in fundamental constants such as as the fine structure constant $\\alpha$ and Newton's gravitational constant $G$. We will not only show the constraints on each single parameter but also the degeneracies among them.  We will consider $3$ experimental configurations: the Planck satellite \\cite{planck}, the combination of Planck with ACT fitted with polarization-sensitive detectors, ACTPol, \\cite{ACTPol} and, finally, the next CMBPol satellite \\cite{CMBPol}.  Recent studies have already fully demonstrated the ability of next generation satellite missions to constrain inflationary parameters \\cite{baumann} and the reionization history \\cite{cmbpolreio} in the framework of the CMBPol concept mission study (see also \\cite{dode}). For this reason we will not consider primordial gravitational waves, more general reionization scenarios or experiments that will mainly probe large angular scale polarization in this paper.  This paper will show that next generation CMB experiments can significantly improve constraints on cosmology and fundamental physics and could produce a detection of neutrino mass. The paper is structured as follows. Section II describes our analysis approach. Section III presents our analysis for improved constraints from the planned ACTPol experiment and for the proposed CMBPol experiment. In Section IV we present our conclusions. ",
        "conclusions": "In this paper we have performed a systematic analysis of the future constraints on several parameters achievable from CMB experiments. Aside from the $5$ parameters of the standard $\\Lambda$-CDM model we have considered new parameters mostly related to quantities which can be probed in a complementary way in the laboratory and/or with astrophysical measurements. In particular we found that the Planck experiment will provide bounds on the sum of the masses $\\Sigma m_{\\nu}$ that could potentially definitively confirm or rule out the Heidelberg-Moscow claim for a detection of an absolute neutrino mass scale. Planck+ACTPol could reach sufficient sensitivity for a robust detection of neutrino mass for an inverted hierarchy, while CMBPol should also be able to detect it for a direct mass hierarchy. The comparison of Planck+ACTPol constraints on baryon density, $N_{eff}$ and $Y_p$ with the complementary bounds from BBN will provide a fundamental test for the whole cosmological scenario. CMBPol could have a very important impact in understanding the epoch of neutrino decoupling. Moreover, the primordial Helium abundance can be constrained with an accuracy equal to that of current astrophysical measurements but with much better control of systematics.  Constraints on fundamental constants can be achieved at a level close to laboratory constraints. Such overlap between cosmology and other fields of physics and astronomy is one of the most interesting aspect of future CMB research.  We should note, however, that our forecasts rely on several technical assumptions. First, we assumed that the theoretical model of the recombination process is accurately known. This is not quite true as corrections to the recombination process are already needed for the Planck experiment (see e.g. \\cite{rubino}). However, this is mainly a computational problem that could be solved by the time of CMBPol launch, expected not before 2015. In addition, we assume that the foreground and beam uncertainties are smaller than the statistical errors.  Nevertheless, the results clearly show the advantage of adding small scale data from the ACT telescope to the Planck satellite data. Adding the former to the latter will improve the constraints by a factor $\\sim2$ in most of the models considered."
    },
    "1005/1005.4934_arXiv.txt": {
        "abstract": "Nearly all globular clusters (GCs) studied to date show evidence for multiple stellar populations, in stark contrast to the conventional view that GCs are a mono-metallic, coeval population of stars.  This generic feature must therefore emerge naturally within massive star cluster formation.  Building on earlier work, we propose a simple physical model for the early evolution (several $10^8$ yr) of GCs. We consider the effects of stellar mass-loss, type II and prompt type Ia supernovae, ram pressure, and accretion from the ambient interstellar medium (ISM) on the development of a young GC's own gas reservoir.  In our model, type II SNe from a first generation of star formation clears the GC of its initial gas reservoir.  Over the next several $10^8$ yr, mass lost from AGB stars and matter accreted from the ambient ISM collect at the center of the GC.  This material must remain quite cool ($T\\sim10^2$K), but does not catastrophically cool on a crossing time because of the high Lyman-Werner flux density in young GCs.  The collection of gas within the GC must compete with ram pressure from the ambient ISM.  After several $10^8$ yr, the Lyman-Werner photon flux density drops by more than three orders of magnitude, allowing molecular hydrogen and then stars to form.  After this second generation of star formation, type II SNe from the second generation and then prompt type Ia SNe associated with the first generation maintain a gas-free GC, thereby ending the cycle of star formation events.  Our model makes clear predictions for the presence or absence of multiple stellar populations within GCs as a function of GC mass and formation environment.  While providing a natural explanation for the approximately equal number of first and second generation stars in GCs, substantial accretion from the ambient ISM may produce fewer chemically peculiar second generation stars than are observed. Analyzing intermediate-age LMC clusters, we find for the first time evidence for a mass threshold of $\\sim10^4\\Msun$ below which LMC clusters appear to be truly coeval.  This threshold mass is consistent with our predictions for the mass at which ram pressure is capable of clearing gas from clusters in the LMC at the present epoch.  Recently, claims have been made that multiple populations within GCs require that GCs form at the center of their own dark matter halos.  We argue that such a scenario is implausible. Observations of the young and intermediate-age clusters in the LMC and M31 will provide strong constraints on our proposed scenario. ",
        "introduction": "\\label{s:intro}  Globular clusters (GCs) have historically been considered coeval, mono-metallic, gravitationally bound collections of stars.  In the past several years, high precision photometric and spectroscopic observations have led to a radical revision of this picture.  High resolution spectra of stars within nearly all GCs studied to date reveal internal spreads in light element abundances such as C, N, Na, O, Mg, and Al, beyond what can be explained by measurement errors \\citep[see review in][]{Gratton04}.  The magnitude of the internal spread varies considerably from cluster to cluster, though there are noticeable trends with cluster mass and orbital properties \\citep{Carretta06, Carretta10b}.  Intriguingly, internal spread in the Fe-peak elements and the type II supernovae product Ca is limited to only the most massive GCs \\citep[$\\omega$Cen, M22, M54, and Terzan 5;][]{Marino09, Carretta09a, Ferraro09, Carretta10a, Carretta10b}. The most massive clusters also show indirect evidence for a large spread in He abundances \\citep{Gratton10}.  Abundance variations have been detected in main sequence stars \\citep{Gratton01}, indicating that the observed variation arises from stars forming out of different material, as opposed to being due to some unknown mixing process, which could only occur along the giant branch.  While much attention has been paid to internal abundance spreads recently, such internal spreads have been known for over 30 years \\citep{Cohen78, Kraft79, Peterson80, Freeman81, Smith82a, Smith83}.  Photometry from the {\\it Hubble Space Telescope (HST)} has demonstrated that many GCs contain multiple sub-giant branches and at least two ($\\omega$Cen and NGC 2808) contain multiple main sequences \\citep[see review in][]{Piotto09}.  Variations in the light element abundances have been conclusively associated with the multiple sequences observed in the color-magnitude diagram \\citep[CMD;][]{Yong08, Marino08, Carretta09b, Milone10}, demonstrating that these two phenomena are intimately linked.  The population with enhanced abundance patterns is more centrally concentrated than the `normal' population in the GCs NGC 1851 \\citep{Zoccali09} and $\\omega$Cen \\citep{Sollima07}, and perhaps many others as well \\citep{Carretta09b}.  In the old Milky Way (MW) GCs the relative numbers of stars with `anomalous' and normal abundances ratios is approximately equal, with little dependence on metallicity \\citep[e.g.,][]{Smith82a, Carretta09b}.  This fact imposes strong constraints on formation scenarios, as we explain in later sections.  The relative numbers of anomalous and normal stars in GCs in other galaxies (e.g., the LMC, M31) is not known, but would provide new insight.  Additional insight has come from the strong observed correlations between various light element abundances within individual GCs.  The most striking is the anti-correlation between Na and O \\citep[e.g.,][]{Kraft93, Ivans99, Carretta09b}.  This correlation arises naturally when material of standard (e.g., solar) abundance ratios is mixed with material that has been processed\\footnote{Throughout the text we will use `processed' and `un-processed' to refer specifically to the nuclear processing of material at the temperatures required to explain the observed light element abundance variations, i.e., at $T>10^7$K.} at temperatures $>10^7$K.  At such temperatures the Na-Ne and CNO cycles are active, with the former producing Na and the latter depleting O.  At similar temperatures the Mg-Al cycle is active, explaining observed correlations between these elements as well.  The discovery of significant Li and F abundances in the second stellar generation \\citep[i.e., associated with stars of high Na and low O abundances;][]{Pasquini05, Smith05} has provided yet another puzzle and clue.  These elements are fragile, especially Li which burns at $T\\gtrsim10^6$K.  The existence of such fragile elements in the atmospheres of second generation stars that also show strong O depletion and Na enhancement suggests that these stars formed out of at least two kinds of material; i.e., from matter exposed to $T>10^7$K, and additional material that was never heated above $T\\sim10^6$K.  The interpretation of Li is however complicated by the possibility that Li may, under special circumstances, actually be produced within AGB stars \\citep[see e.g.,][]{Ventura10}.  Multiple stellar populations have also been detected in intermediate-age ($\\sim1$ Gyr) and old LMC clusters \\citep{Mackey08, Goudfroij09, Milone09, Mucciarelli09}, and possibly in the old GCs within the Fornax dSph galaxy \\citep{Letarte06}.  These observations indicate that the multiple stellar population phenomenon is not specific to the MW.  Observations of LMC clusters are of course hampered by the much larger distance modulus to the LMC.  Despite this limitation, the intermediate-age clusters offer a new window into the internal age spreads because the main sequence turn-off point at these ages is a strong function of time.  Small age differences are therefore readily noticeable in the CMD.  Analysis of {\\it HST}-based CMDs have shown that the spread observed in the main sequence turn-off of intermediate-age LMC clusters can be explained with an internal age spread of a few $10^8$ yr.  The only class of star clusters known {\\it not} to contain multiple stellar populations are the open clusters in the MW \\citep{deSilva09, Martell09}, which have typical masses of a few thousand solar masses, although some, such as NGC 7789, have masses of $\\sim10^4\\Msun$.  These observations have led to the unavoidable conclusion that the majority of GCs studied to date harbor multiple stellar populations. For all but the most massive ones, GCs are still considered chemically homogeneous in Fe-peak and elements arising primarily from SNe type II such as Ca.  The light element variations have been detected in both young and old clusters, both metal-poor and metal-rich \\citep{Martell09}, and are noticeably absent in the open clusters.  From these observations the following timeline in the early evolution of GCs has emerged.  Within GCs, a first generation of stars form. Type II SNe then remove any remaining gas from the GC.  After the epoch of type II SNe, mass from evolved stars is cycled through temperatures of $T>10^7$K and then is returned to the gaseous reservoir of the GC.  After a few $10^8$ yr, a second generation of stars forms from a mix of processed and un-processed material.  Star formation permanently ceases after the formation of the second generation.  This process does not occur in open clusters, nor in the field.  This basic timeline has, in one form or another, been discussed by many authors \\citep[e.g.,][]{Cottrell81, Smith87, Carretta10c}.  Any theory of GC formation must be embedded into our broader cosmological theories of structure formation.  Many lines of evidence suggest that galaxy formation is an hierarchical process where dark matter halos serve as sites for assembling baryons and converting baryons into stars.  One of the inevitable predictions of this bottom-up scenario is that MW GCs likely formed in environments very different from the $\\sim200\\kms$ dark matter dominated halo where they now reside.  One possibility, which we discuss in $\\S$\\ref{s:dm}, is that GCs form in the centers of small dark matter halos.  This scenario would imply that isolated GC are embedded in extended dark matter halos.  This paper emphasizes another possibility: GCs form within small gas-rich dwarf galaxies --- the building blocks of the present MW.  We emphasize that the MW is an evolving galactic system so that estimates of ram pressure and gaseous accretion must consider the likely environment in which a GC formed at $z\\sim2-10$, rather than its current location in the MW today.  We have outlined above only the most basic sketch of what must occur to explain the observations.  In the next section we critically assess previous, more detailed scenarios for the development of multiple stellar populations within GCs.  Following this assessment, we describe our own model for the early evolution of GC stellar populations that includes several novel ingredients, and is, at least in certain respects, more plausible than other scenarios.  We also present an analysis of intermediate-age clusters in the LMC that provides confirmation of a key aspect of our model.  We conclude by commenting on various observations that may shed new light on this exciting observational puzzle. ",
        "conclusions": "\\label{s:res}  In this work we have presented a comprehensive model for the early (several $10^8$ yr) evolution of massive star clusters.  Our model considers the importance of type II and prompt type Ia SNe, accretion onto the GC from the ambient ISM, the ability of a GC to retain its internal gas supply in the face of ram pressure, and the effect of the Lyman-Werner photon flux density on the ability of the young GC gas to form molecules and, ultimately, stars.  This model definitively addresses two of the three major issues in early GC evolution identified in $\\S$\\ref{s:prev}: how GCs can retain a gaseous reservoir in the face of ram pressure, and which stars are responsible for processing material to $T>10^7$K.  From consideration of the formation environments of both old MW GCs and intermediate-age LMC clusters, we have shown that ram pressure stripping naturally explains the observed bifurcation between clusters that do and do not show evidence for multiple stellar populations.  Based on several independent arguments, including the similar $\\sim10^8$ yr timescale of many physical mechanisms and the lack of internal spread in Fe and Ca in most clusters, we strongly favor massive AGB stars as the source of the processed material.  The final open issue is in understanding how a second generation can form with a current total mass comparable to the first generation, at least for old MW GCs where data are available.  We have considered accretion from the ambient ISM as a viable mechanism to provide copious amounts of gas to the GC.  However, our solution to this last issue currently has little direct empirical verification, and may in fact have trouble reproducing the observed correlations in abundance ratios, as discussed below.  We now discuss a variety of testable implications of our proposed solution, and comment briefly on several open issues.  As a consequence of significant accretion from the ambient ISM, we expect several trends with GC mass (see e.g., Figure \\ref{fig:facc}). For example, as the fraction of pristine gaseous material increases, we expect a shorter Na-O correlation, a greater abundance of Li and F, and in general we expect the stars with anomalous abundances to be less anomalous.  The sign of the trend with mass depends on the dominant accretion process, which in turn depends on the formation environment.  Bondi accretion scales as $M^2$ while the mass provided by AGB winds scales with $M$ so the pristine fraction increases with mass.  ISM sweeping scales as $R^2$ which is only weakly correlated with $M$ at the present epoch, and so if this process is dominant we would expect the pristine fraction to decrease with mass.  Observations of young and intermediate-age clusters will allow identification of the relevant mechanism because the formation environment of young clusters is not so dissimilar from their present environment.  The old GCs within dwarf spheroidals and dwarf irregulars may also shed light on this issue.  The GCs within dwarf galaxies almost certainly formed where they are now observed since the stellar accretion/merger rate onto dwarfs is expected to be very low in a $\\Lambda$CDM cosmology, and so the present day conditions of the host dwarf galaxies cannot be very different from the formation environment of the GCs.  \\citet{Carretta06} and \\citet{Carretta10c} have investigated the extent of the Na-O correlation as a function of global parameters and have found evidence that the correlation is more extended in higher mass GCs, and for GCs with more extended orbits.  This result is consistent with ISM sweeping being the dominant accretion process, but we caution that the data show large scatter, and the present GC mass is poorly correlated with its mass at formation due to orbit-dependent mass-loss effects.  In general, we expect that the extent of the Na-O correlation should depend on the formation environment at fixed GC mass, since the formation environment determines the amount of accreted material.  Another effect might be the mass-dependent ability of a GC to retain the winds from AGB stars.  In any event, it is clear that such trends hold the promise of isolating which physical process dominates the accretion of pristine material onto GCs.  If accretion from the ambient ISM was important, then it is somewhat puzzling why the Fe and Ca abundances are so uniform between the first and second generations.  The only plausible explanation is that at early times, during the formation of the ancient MW GCs, the spread in Fe and Ca abundances within the MW progenitor system was quite small. A clear prediction of the accretion scenario is that the intermediate-age LMC clusters showing multiple populations should show a much larger internal spread in Fe, Ca, etc. abundances, on the order of the spread in abundances of these elements within the LMC as a whole.  The MW bulge GC Terzan 5 may provide additional insight. \\citet{Ferraro09} has convincingly shown that this GC has a split horizontal branch (HB) with the more luminous branch having a much higher Fe abundance ([Fe/H]$\\sim+0.3$) than the less luminous branch ([Fe/H]$\\sim-0.2$).  \\citet{DAntona10} argue that the split HB is consistent with an internal age spread of several $10^8$ yr and enhanced He in the metal-enriched branch.  Since this cluster is in the MW bulge and Fe-rich, it probably formed in the bulge.  D'Antona et al. suggest that the second stellar generation may have acquired its high Fe abundance via accretion from the ambient, Fe-rich ISM. Self-enrichment is unlikely because the required Fe mass to produce the observed high Fe abundance would require so many type II SNe that the cluster would easily become unbound, unless Terzan 5 was significantly more massive in the past.  Terzan 5 may therefore be an ideal cluster to look for further evidence for the importance of ambient ISM accretion in the formation of multiple stellar populations, e.g., by observing the Li and F abundances in this cluster.  If the proposed scenario for Terzan 5 is generic then we might expect to find significantly separated Fe abundances within the other bulge metal-rich GCs as well.  In the present work we have not attempted to make specific predictions for the extent of correlations amongst various elemental abundances. Such predictions would provide a powerful constraint on the model. Qualitative considerations suggest that substantial accretion from the ISM may yield a very short Na-O anti-correlation, owing to the ISM having abundance ratios similar to the first generation stars. Unfortunately, uncertainties in the AGB yields greatly complicate any attempted comparison to observed abundances.  A fruitful avenue for future work will consist of a detailed comparison between predicted and observed abundance correlations in light of the uncertain AGB yields \\citep[see][for an initial attempt in this direction]{DErcole10b}.  The most massive GCs, including $\\omega$Cen, NGC 2808, M22, and M54, deserve special mention.  Each of these clusters show multiple, {\\it distinct} sequences in the CMD, and the former two are unique in that they display multiple main sequences.  The CMD morphology of these clusters suggests very high He abundances in the second (and third) stellar generations.  The discrete sequences and high He abundances argue against significant dilution from ambient ISM accretion.  This can be accommodated in our model if ISM sweeping is the dominant accretion mechanism because this mechanism becomes increasingly less important at high GC masses (see the bottom panels of Figure \\ref{fig:facc}).  Although we are then left with the original problem of explaining how so many second generation stars could form from the mass lost by the first generation.  We speculate that perhaps the most massive GCs formed at the centers of their own dark matter halos, and contained even more stars in their past.  The precursors of these systems could be nuclear star clusters, which are common in low-mass galaxies.  Regardless of these details, we stress caution when attempting to incorporate the massive GCs into any framework for the early evolution of GCs since many of their observational characteristics differ qualitatively from the lower mass GCs.  An exciting direction for future observational work is characterizing the incidence of multiple populations within the numerous young and intermediate-age clusters in M31.  These clusters span a range in mass from $10^3\\Msun$ to $10^5\\Msun$ \\citep{Caldwell09}.  Based on ram pressure arguments, we expect the higher mass clusters to show evidence for multiple populations, but not the lower mass clusters. Assuming $n=1\\cm$ and $V=250\\kms$ for disk of M31 \\citep{Widrow03}, we predict that the critical mass will be $\\approx10^{4.5}\\Msun$. However, as stressed in previous sections, the relevant velocity is the relative velocity between the ambient ISM and the GC.  Since many of the young M31 GCs are orbiting within the disk of M31, their relative velocity will be considerably less than the circular velocity, and so the critical mass may be considerably lower. Confrontation of our model predictions with the properties of the M31 GCs must take these details into account.  Since these M31 clusters have not experienced significant dynamical evolution, we may expect stronger correlations between the extent of any abundance anomalies and cluster mass.  For the youngest clusters, with ages of several $10^8$ yr, we might even hope to {\\it directly observe} the formation of the second generation.  Such an observation would provide definitive proof that GCs need not form within dark matter halos to produce multiple stellar generations, and would also demonstrate that AGB stars are the polluters, since the AGB polluter scenario is expected to produce a second generation on a timescale of several $10^8$ yr.  In the past several years it has become clear that star clusters, once thought to be simple systems, in fact show an internal complexity that increases with increasing mass.  Observations indicate that the lowest mass systems, i.e., the open clusters, appear to be truly coeval and mono-metallic.  At higher masses one observes age spreads of several $10^8$ yr and internal variation in the light elements.  At still higher masses ($\\gtrsim10^6\\Msun$) it appears that the cluster potential well is deep enough to retain SNe ejecta and hence self-enrich, perhaps because these massive clusters form at the center of their own dark matter halos.  This last phenomenon is observed not only in the most massive MW GCs, but also appears to be present in the most massive GCs within other galaxies as well \\citep{Strader08, Bailin09}.  At the very least, it is now abundantly clear that star clusters are not simple systems."
    },
    "1005/1005.1720_arXiv.txt": {
        "abstract": "In this letter, we study in detail the evolution of the post-flare loops on 2005 January 15 that occurred between two consecutive solar eruption events, both of which generated a fast halo CME and a major flare. The post-flare loop system, formed after the first CME/flare eruption, evolved rapidly, as manifested by the unusual accelerating rise motion of the loops. Through nonlinear force-free field (NLFFF) models, we obtain the magnetic structure over the active region. It clearly shows that the flux rope below the loops also kept rising accompanied with increasing twist and length. Finally, the post-flare magnetic configuration evolved to a state that resulted in the second CME/flare eruption. This is an event in which the post-flare loops can re-flare in a short period of $\\sim$16 hr following the first CME/flare eruption. The observed re-flaring at the same location is likely driven by the rapid evolution of the flux rope caused by the magnetic flux emergence and the rotation of the sunspot. This observation provides valuable information on CME/flare models and their prediction. ",
        "introduction": "The developed systems of loops, named post-flare loops, usually occur in the decay phase of long duration events \\citep[e.g.,][]{bruzek64,kahler77,Harra98}. It is generally believed that magnetic reconnection plays an important role in the formation and evolution of the post-flare loops. The rising of post-flare loops in the corona, as well as the separation of flare ribbons at the footpoints, can be well explained by the classical reconnection model, i.e., the CSHKP model \\citep{carmichael64,sturrock66,hirayama74,kopp76}. In this model, the observed motions are caused by the systematic ascending of the reconnection site in the corona \\citep[e.g.,][]{cheng10a}. The post-flare loops comprise of a system of loops heated and formed consecutively through the ongoing reconnection process. Based on the CSHKP model, the post-flare loops rise with a gradually decreasing speed. However, \\citet{svestka96} found some exceptional cases, in which the post-flare loops rose with a constant speed  for a long period of time, and was explained through the combination of two entirely different processes: the initial development of the magnetic reconnection proposed by \\citet{kopp76} and the subsequent expanding motion into the corona. It is generally believed that the post-flare loops evolve without further energy input after their formation. Therefore, they usually fade away into the background after a certain time. To our knowledge, there is no report on the re-activation of a post-flare loop system in previous studies.   Coronal mass ejections (CMEs) are another kind of large-scale solar activity that may involve the magnetic field lines that twisted each other, named flux rope, in their core, as revealed by many coronagraph observations \\citep[e.g.,][]{pick06}. The role of flux ropes in the CME eruptions has also been studied extensively through numerical simulation. \\citet{amari00} simulated the evolution of a flux rope and found that it plays a crucial role in the process of the CME eruption. The flux rope was also used by \\citet{fan01} and \\citet{fan03,fan04} to investigate the dynamic evolution of the coronal magnetic field in response to the emergence of twisted magnetic structures. Moreover, \\citet{torok03,torok05} and \\citet{Kliem06} studied the instability of the flux rope and proposed that the CME eruption can be initiated through the kink and/or torus instability. \\citet{cheng10b} suggested that the flux rope may provide a magnetic structure favorable for the eruption and regarded the formation period of the flux rope prior to the eruption as the build-up phase of CMEs \\citep[see also][]{su09,Aulanier10}. \\citet{schrijver08} and \\citet{schrijver09} showed that the emergence of current-carrying flux ropes can lead to the occurrence of a series of major flares within a few hours in the same active region. Further, \\citet{canou09} constructed a flux rope structure above magnetic polarity inversion line (PIL) of an emerging sunspot by a nonlinear force-free field (NLFFF) model using the THEMIS vector magnetogram as the input. Recently, \\citet{guo10} also obtained a flux rope structure through NLFFF modeling and found that the flux rope coincided with the active region filament. However, up to now, there is still no study on the time evolution of extrapolated flux ropes.   In this letter, we report the re-flaring of a post-flare loop system occurred on 2005 January 15, and further study the temporal development of associated flux rope. In section 2, we present the observations and data analysis method. Our results are shown in section 3 followed by summary and discussions in section 4.  \\begin{figure*}[!ht] \\begin{centering} \\epsscale{0.9} \\plotone{cmeflare.eps} \\caption{(a--b) Running difference images of CMEs 1 and 2 observed by LASCO. (c) EIT 195 {\\AA} image of the CME-associated Flare 1. (d--f) Three images showing the post-flare loops after the peak time of flare 1. Slices 1 and 2 in panel (f) are used to trace the time evolution of the post-flare loop system. (g) EIT 195 {\\AA} image of the CME-associated flare 2. (h) Image showing the post-flare loops of flare 2. \\label{cmeflare}}  (Animation of the post-flare loops is available in the online journal) \\end{centering} \\end{figure*} ",
        "conclusions": "In this study, we report the full evolution of a post-flare loop system, characterized by the initial rising and expanding motion and the final re-activation as another CME. We find that the rising and expanding of the post-flare loops were essentially due to the continual increase of the magnetic flux, which enhanced the magnetic pressure underneath the post-flare loop system. On the other hand, we speculate that the first CME/flare eruption may lead to a weakening of the constraining tension force of the overlying field to the post-flare loop system. We further study the development of the associated flux rope by the NLFFF extrapolation, and find that the twist and length of the flux rope were increased, which may be due to the continual flux emergence and/or the rotation of the sunspot. In fact, magnetic flux emergence \\citep{fan01,fan03,fan04} and sunspots rotation caused by photospheric plasma flows \\citep{amari00,torok03,torok05} are among the most plausible ways that have been invoked in numerical simulations to increase the magnetic non-potentiality leading to the final eruption.   Through a quasi-linear force free method \\citep{wang01}, \\citet{zhao08} studied the magnetic topology skeleton of this active region. They found the rising of the magnetic null and its spine, which is consistent with the rising of the extrapolated flux rope and the strong twist of the flux rope as shown here. This strongly twisted flux rope provides a source for the shear of the post-flare loop system that enhances the non-potentiality of associated magnetic configuration. \\citet{jing09} studied the temporal evolution of the magnetic free energy for this active region from 21:00 to 23:00 UT on January 15 and found that it tends to increase prior to the onset of flare 2. The accumulation of the free energy was mainly caused by the emergence of magnetic flux and the rotation of the sunspots, which continually transfers energy below the photosphere into the corona \\citep{kurokawa02,cheng10b}. Therefore, the post-flare configuration may still accumulate magnetic free energy in a short time that is responsible for the second eruption.  It is worthy noting that the homologous flares can take place in the same active region with similar morphologies and pattern of development \\citep[e.g.,][]{choe00}, and the successive flares usually occur in the magnetic loops that are closely related but not in exactly the same loops \\citep[e.g.,][]{liu09}. Whereas, the most significant difference, of the events in this studies from the general homologous and successive flares, is that the second CME/flare came directly from the post-flare loops of the first eruption but not the loops structures in the similar or other related locations.    In conclusion, it is commonly viewed that a post-flare loop system evolves relatively uneventful after their formation following a CME/flare eruption. Except for some dynamic flows along the post-flare loops \\citep[e.g.,][]{1997SoPh..175..511B,1999SoPh..190..153Q}, the magnetic structure of the loops is closer to potential than that before the eruption, thus unable to re-erupt in a short time after a major eruption. The event studied here provides an exceptional case and may break such a point of view. It shows that a post-flare magnetic configuration can evolve quickly to a state that is favorable for another major CME/flare eruption, as long as there is continual emergence of flux rope and/or rotation of sunspots. Post-flare loops can re-flare in a short time. Models and prediction of solar flares (CMEs) should take into account such an observational fact."
    },
    "1005/1005.5335_arXiv.txt": {
        "abstract": "{ Dynamos in the Sun and other bodies tend to produce magnetic fields that possess magnetic helicity of opposite sign at large and small scales, respectively. The build-up of magnetic helicity at small scales provides an important saturation mechanism. } { In order to understand the nature of the solar dynamo we need to understand the details of the saturation mechanism in spherical geometry. In particular, we want to understand the effects of magnetic helicity fluxes from turbulence and meridional circulation. } { We consider a model with just radial shear confined to a thin layer (tachocline) at the bottom of the convection zone. The kinetic $\\alpha$ owing to helical turbulence is assumed to be localized in a region above the convection zone. The dynamical quenching formalism is used to describe the build-up of mean magnetic helicity in the model, which results in a magnetic $\\alpha$ effect that feeds back on the kinetic $\\alpha$ effect. In some cases we compare with results obtained using a simple algebraic $\\alpha$ quenching formula. } { In agreement with earlier findings, the magnetic $\\alpha$ effect in the dynamical $\\alpha$ quenching formalism has the opposite sign compared with the kinetic $\\alpha$ effect and leads to a catastrophic decrease of the saturation field strength with increasing magnetic Reynolds numbers. However, at high latitudes this quenching effect can lead to secondary dynamo waves that propagate poleward due to the opposite sign of $\\alpha$. Magnetic helicity fluxes both from turbulent mixing and from meridional circulation alleviate catastrophic quenching. } {} ",
        "introduction": "The solar dynamo models developed so far and which agree with solar magnetic field observations tend to solve the $\\alpha\\Omega$ mean field dynamo equations. The turbulent $\\alpha$-effect first proposed by  Parker (1955) is believed to be generated due to helical turbulence in the convection zone of the Sun.  Since $\\alpha$ is generated due to quadratic correlations of the small-scale turbulence we need a closure in order to complete the set of mean field equations, e.g., the first order smoothing approximation (FOSA), and express the mean electromotive force in terms of the mean magnetic fields. This turbulent $\\alpha$ encounters a critical problem when the energy of the mean field becomes comparable to the equipartition energy of the turbulence in the convection  zone and hence it becomes increasingly difficult for the helical turbulence to twist rising blobs of magnetic field. The solar dynamo modellers have traditionally used  what is referred to as algebraic alpha quenching to mimic this phenomena.  This involves replacing $\\alpha$ by $\\alpha_0/(1+\\ob{B}^2/B_{\\rm eq}^2)$,  an expression used since Jepps (1975), or by $\\alpha_0/(1+R_{\\rm m} \\ob{B}^2/B_{\\rm eq}^2)$, where $\\alpha_0$ is the unquenched value and $R_{\\rm m}$ is the magnetic  Reynolds number, $\\ob{B}$ is the mean magnetic field and $B_{\\rm eq}$ is the equipartition magnetic field. The latter expression has been discussed since the early work of Vainshtein \\& Cattaneo (1992). The $R_{\\rm m}$ in the denominator comes from  the fact that the  small-scale fluctuating magnetic field reaches equipartition long  before the mean magnetic field does. This has been supported by several numerical  experiments to determine the saturation behaviour of $\\alpha$ (e.g. Cattaneo \\& Hughes 1996, Ossendrijver et al. 2002). Given the large magnetic Reynolds numbers of Astronomical objects, such phenomena is referred to as catastrophic quenching.  After the discovery of the layer of strong radial shear (called the tachocline by Spiegel \\& Zahn 1992) at the bottom of the solar convection zone, Parker (1993) proposed a new class of solar dynamo models called the interface dynamo. In these models the shear is confined to a narrow overshoot layer just beneath the convection zone, also the region of $\\alpha$ effect. The dynamo wave propagates in a direction given by the Parker--Yoshimura rule at the interface between the two layers defined by a steep gradient in the turbulent diffusivity. The toroidal field produced due to stretching by the shear is much stronger than the poloidal field and remains confined in the overshoot layer, away from the region where the $\\alpha$ effect operates.  It may be noted that the interface dynamo model may have serious problems when solar-like rotation with positive latitudinal shear is included (Markiel \\& Thomas 1999). Similarly, in the Babcock-Leighton class of flux transport models (Choudhuri et al.\\ 1995; Durney 1995) the toroidal and the poloidal fields  are produced in two different layers. Unlike in the interface dynamo models, the coupling between the two layers is mediated both by diffusion and the conveyer belt mechanism of the meridional circulation.  It has been proposed that in interface and Babcock-Leighton type dynamos,  the $\\alpha$ effect is not catastrophically quenched at high $\\Rm$ because the strength of the toroidal field is very weak in the region of finite turbulent $\\alpha$ (e.g. Tobias, 1996; Charbonneau, 2005). However, according to our knowledge, not much has been done to study the variation of the amplitude of the saturation magnetic field with the magnetic Reynolds  number for these classes of  $\\alpha\\Omega$ dynamos.  Zhang et al (2006) made an attempt to reproduce the surface  observations of current helicity in the Sun using a 2D mean field  dynamo model in spherical coordinates coupled with the dynamical  quenching equation. In a separate paper (Chatterjee, Brandenburg \\&  Guerrero, 2010) we have demonstrated that interface dynamo models are  also subject  to catastrophic quenching.  It has been identified a decade ago that the small-scale magnetic helicity generated due to the dynamo action back reacts on the helical turbulence and quenches the dynamo (Blackman \\& Field, 2000; Kleeorin et al. 2000). It has now been shown that this mechanism reduces the saturation amplitude of the magnetic field ($B_{sat}$) with increasing magnetic Reynolds number ($\\Rm$). Nevertheless this constraint may be lifted if the system is able to get rid of small scale helicity through several ways like open boundaries, advective, diffusive and shear driven fluxes (Shukurov et al. 2006, Zhang et al. 2006, Sur et al. 2007, K\\\"apyl\\\"a et al. 2008, Brandenburg et al. 2009, Guerrero et al. 2010). Even though the helicity constraint in direct numerical simulations (DNS) of dynamos with strong shear have been clearly identified, the results can be matched with mean field models having a weaker algebraic quenching than $\\alpha^2$ dynamos (Brandenburg et al.\\ 2001). It is possible to include this process in mean-field dynamo models through an equation describing the evolution of the small scale current helicity. We shall refer to this equation as the dynamical quenching mechanism.  In this paper we perform a series of calculations with mean field $\\alpha\\Omega$ models in spherical geometry along with a dynamical equation for the evolution of $\\alpha$ for magnetic Reynolds numbers in the range $1\\le R_{\\rm m}\\le 2\\times10^5$. An important feature of the calculation is that the region of strong narrow shear is separated from the region of helical turbulence. This paper in addition to providing detailed results not mentioned in Chatterjee, Brandenburg \\& Guerrero (2010), is also aimed at studying somewhat more complicated models including meridional circulation. The role of diffusive helicity fluxes modelled into the dynamical quenching equation by using a Fickian diffusion term is also discussed for various models. It may be mentioned that helicity fluxes across an equator can indeed be modelled by such a diffusion term as shown by Mitra et al. (2010). In \\S2 we discuss the features of the $\\alpha\\Omega$ model used, and the formulation of dynamical $\\alpha$ quenching. The results are highlighted in \\S3 and  conclusions are drawn in \\S4. ",
        "conclusions": "We have performed calculations for $\\alpha\\Omega$ dynamos in a spherical shell for spatially segregated $\\alpha$ and $\\Omega$ source regions. The two classes of models we have studied resemble the Parker's interface dynamo and the Babcock-Leighton dynamo.  In agreement with earlier work, it is not possible to escape catastrophic quenching by merely separating the regions of shear and $\\alpha$-effect. The saturation value of magnetic energy decreases as $\\sim R_{\\rm m}^{-1}$ for both dynamical quenching and the algebraic quenching with $g_{\\alpha}=\\Rm$ for the simple two layer model without meridional circulation (Fig.~\\ref{fig:satrm}). However we find that a richer dynamical behaviour emerges for the cases with dynamical $\\alpha$ effect, in terms of parity fluctuations and appearance of `secondary' dynamos (Fig.~8, 9). We do not see evidence for chaotic behaviour in the time series of magnetic energy since the dynamo period and the saturation energy remains fairly constant. However this may not be the case in presence of diffusive helicity fluxes which introduce further complexity to the system. Addition of diffusive helicity fluxes relaxes the catastrophic $\\Rm^{-1}$ dependence of the saturation magnetic energy (Fig.~10a, 12). An interesting `side-effect' of diffusive helicity fluxes is the appearance of poleward propagating secondary dynamos. However, because of the lack of scale separation between the mean field and the forcing scale of the helical turbulence we refrain from interpreting this in terms of the poleward migration seen in the Sun. It remains to explore the role of the solar wind, coronal mass ejections which might help in throwing out the small scale helicity from the Sun and thus alleviate catastrophic quenching. The effects of Vishniac \\& Cho fluxes have been investigated and were found to be of secondary importance compared to diffusive helicity fluxes for $\\alpha\\Omega$ mean field dynamos (Guerrero, Chatterjee \\& Brandenburg 2010).  When both the meridional circulation and the diffusive helicity fluxes are artificially shut off in the helicity evolution equation, the dynamo fails to reach significant saturation values, as expected (Fig.~12). It is interesting that the Babcock-Leighton dynamos, where $\\alpha$ is concentrated only in a narrow layer at the surface, also produce considerable helicity inside the convection zone when the dynamical quenching (Eq.~\\ref{eq:alphaeq}) is employed (Fig.~13, 14).  We have to be cautious about using dynamical quenching equation for dynamo numbers not very large compared to the critical dynamo number. For highly supercritical $\\alpha$, the behaviour of the system begins to be governed by $\\alpha_{\\rm M}$. We would expect that the magnetic field should affect all the turbulent coefficients including both $\\alpha$ and $\\eta$. However for this analysis we have not included an equation for the variation for $\\eta_{\\rm t}$. This is justified for the simple two layer model with a lower $\\eta_{\\rm t}$ in the region of production of strong toroidal fields and a higher $\\eta_{\\rm t}$ in the region of weaker poloidal fields. It may also be noted that by quenching the diffusivity inversely with the magnetic energy in a nonlinear dynamo model, Tobias (1996) was able to produce a bonafide interface model where the magnetic field was restricted to a thin layer at an interface between a layer of shear and cyclonic turbulence. However none of the previous interface models have used the dynamical quenching equation.  Unfortunately the direct numerical simulations have not yet reached the modest Reynolds numbers used in this paper ($\\sim 10^4$) which are still much lower than the astrophysical dynamos. To verify if the equation for dynamical quenching works in the same way as in $\\alpha^2$ dynamos, we need to embark upon systematic comparisons between DNS with shear and convection and mean field modelling for $\\alpha\\Omega$ dynamos."
    },
    "1005/1005.0383_arXiv.txt": {
        "abstract": "{Significant progress has been made in the last years in the understanding of the jet formation mechanism through a combination of numerical simulations and analytical MHD models for outflows characterized by the symmetry of self-similarity. Analytical radially self-similar models successfully describe disk-winds, but need several improvements. In a previous article we introduced models of truncated jets from disks, i.e. evolved in time numerical simulations based on a radially self-similar MHD solution, but including the effects of a finite radius of the jet-emitting disk and thus the outflow.} {These models need now to be compared with available observational data. A direct comparison of the results of combined analytical theoretical models and numerical simulations with observations has not been performed as yet. This is our main goal.} {In order to compare our models with observed jet widths inferred from recent optical images taken with the Hubble Space Telescope (HST) and ground-based adaptive optics (AO) observations, we use a new set of tools to create emission maps in different forbidden lines, from which we determine the jet width as the full-width half-maximum of the emission.} {It is shown that the untruncated analytical disk outflow solution considered here cannot fit the small jet widths inferred by observations of several jets. Furthermore, various truncated disk-wind models are examined, whose extracted jet widths range from higher to lower values compared to the observations. Thus, we can fit the observed range of jet widths by tuning our models.} {We conclude that truncation is necessary to reproduce the observed jet widths and our simulations limit the possible range of truncation radii. We infer that the truncation radius, which is the radius on the disk mid-plane where the jet-emitting disk switches to a standard disk, must be between around 0.1 up to about 1 AU in the observed sample for the considered disk-wind solution. One disk-wind simulation with an {{\\em inner}} truncation radius at about 0.11 AU also shows potential for reproducing the observations, but a parameter study is needed.} ",
        "introduction": "Astrophysical jets and disks \\citep{Liv09} seem to be inter-related, notably in young stellar objects (YSOs), where jet signatures are well correlated with the infrared excess and accretion rate of the circumstellar disk \\citep{CES90, HEP04}. Disks provide the plasma which is outflowing in the jets, while jets in turn provide the disk with the needed angular momentum removal so that accretion onto the protostellar object takes place \\citep{Har09}. On the theoretical front, the most widely accepted description of this accretion-ejection phenomenon \\citep{Fer07} is based on the interaction of a large scale magnetic field with an accretion disk around the central object. Then, plasma is channeled and magneto-centrifugally accelerated along the open magnetic field lines threading the accretion disk, as first described in \\citet{BlP82}. Several works have extended this study either by semi-analytic models using radially self-similar solutions of the full magnetohydrodynamics (MHD) equations with the disk treated as a boundary condition \\citep{VlT98}, by self-consistently treating the disk-jet system semi-analytically \\citep[e.g.][]{Fer97,CaF00}, or, by self-consistently treating numerically the disk-jet system \\citep[e.g.][]{ZFR07, TFM09}.  The original \\citet{BlP82} model, however, has serious limitations for a needed meaningful comparison of its predictions with observations. First, singularities exist at the jet axis, the outflow is not asymptotically super-fast, and most importantly, an intrinsic scale in the disk is lacking with the result that the jet formally extends to radial infinity, to mention just a few. First, the singularity at the axis can be easily taken care of by numerical simulations extending the analytical solutions close to this symmetry axis \\citep[][GVT06 hereafter]{GVT06}. Next, the outflow speed at large distances may be tuned to cross the corresponding limiting characteristic, with the result that the terminal wind solution is causally disconnected from the disk and hence perturbations downstream of the super-fast transition (as modified by self-similarity) cannot affect the whole structure of the steady disk-wind outflow \\citep[][V00 hereafter]{VTS00}, a state which has also been shown to be structurally stable \\citep[][M08 hereafter]{MTV08}. The next step of introducing a scale in the disk has been done in a previous paper \\citep[][paper I hereafter]{STV08}, wherein we presented numerical simulations of truncated flows whose initial conditions are based on analytical self-similar models.  \\begin{table*}[!htb] \\caption{List of numerical science models} \\label{tbl_models} \\centering \\begin{tabular}{l l l l l} \\hline\\hline Name & Domain $[ R_0$ x $R_0]$ & Resolution & Description & $R_{\\rm trunc}$ $[ R_0 ]$  \\\\ \\hline model SC1a & [0,50] $\\times$ [6,100]  & 200 $\\times$ 400 & $\\alpha_{\\rm trunc} = 0.4$, external analytical solution $\\lambda_1 = 10^3$, $\\lambda_2 = 10^{-3}$ & 5.375 \\\\ model SC1b & [0,50] $\\times$ [6,100]  & 200 $\\times$ 400 & $\\alpha_{\\rm trunc} = 0.2$, external analytical solution $\\lambda_1 = 10^3$, $\\lambda_2 = 10^{-3}$ & 5.125 \\\\ model SC1c & [0,50] $\\times$ [6,100]  & 200 $\\times$ 400 & $\\alpha_{\\rm trunc} = 0.1$, external analytical solution $\\lambda_1 = 10^3$, $\\lambda_2 = 10^{-3}$ & 4.875 \\\\ model SC1d & [0,50] $\\times$ [6,100]  & 200 $\\times$ 400 & $\\alpha_{\\rm trunc} = 0.01$, external analytical solution $\\lambda_1 = 10^3$, $\\lambda_2 = 10^{-3}$ & 3.625 \\\\ model SC1e & [0,50] $\\times$ [6,100]  & 200 $\\times$ 400 & $\\alpha_{\\rm trunc} = 0.001$, external analytical solution $\\lambda_1 = 10^3$, $\\lambda_2 = 10^{-3}$ & 2.625 \\\\ model SC2     & [0,50] $\\times$ [6,100]  & 200 $\\times$ 400 & $\\alpha_{\\rm trunc} = 0.4$, external analytical solution $\\lambda_1 = 100$, $\\lambda_2 = 0.1$ & 5.375 \\\\ model SC3     & [0,50] $\\times$ [6,100]  & 200 $\\times$ 400 & same as model SC2, but solutions are swapped & 5.375 \\\\ model SC4     & [0,50] $\\times$ [6,100]  & 200 $\\times$ 400 & $\\alpha_{\\rm trunc} = 0.4$, external analytical solution $\\lambda_1 = 1$, $\\lambda_2 = 0.1$ & 5.375 \\\\ model SC5     & [0,50] $\\times$ [6,100]  & 200 $\\times$ 400 & same as model SC4, but solutions are swapped & 5.375 \\\\ \\hline \\hline model SC1f & [0,50] $\\times$ [6,100]  & 200 $\\times$ 400 & $\\alpha_{\\rm trunc} = 0.0005$, external analytical solution $\\lambda_1 = 10^3$, $\\lambda_2 = 10^{-3}$ & 2.375 \\\\ model SC1g & [0,10] $\\times$ [6,20]  & 200 $\\times$ 400 & $\\alpha_{\\rm trunc} = 1\\times10^{-5}$, external analytical solution $\\lambda_1 = 10^3$, $\\lambda_2 = 10^{-3}$ & 0.575 \\\\ \\hline \\end{tabular} \\end{table*}  In order to test our truncated models, we will now apply our simulations to observations. In recent years, many NIR and optical data have become available exploring the morphology and kinematics of the jet launching region \\citep[e.g.][and references therein]{DCL00,RDB07,Dou08}. Hubble Space Telescope (HST) and adaptive optics (AO) observations give access to the innermost regions of the wind, where the acceleration and collimation occurs \\citep{RMD96,DCL00,WRB02,HEP04}. Because YSO jets emit in a number of atomic (and molecular) lines, we used a set of tools described in \\citet{G??10} to create emission maps in different forbidden lines which were used by other authors to extract the jet width from images. The observed jet widths will be compared with those extracted from our synthetic images. A similar study has been done by \\citet{CFR99}, \\citet{GCF01} and \\citet{DCF04} using a different set of semi-analytical self-similar disk-jet solutions from \\citet{Fer97} and \\citet{CaF00}. Observed jet widths could be reproduced by manually truncating the solutions inside 0.07 AU and outside 1 AU, but the modification of flow streamlines induced by truncation was ignored.  The remainder of the paper is organized as follows: we briefly review the initial setup of the numerical simulations in Sect. \\ref{sec_num_models} and describe our procedure for the comparison with observations in Sect. \\ref{sec_obs}. The results of our studies are presented in Sects. \\ref{sec_untrunc} - \\ref{sec_inner}. Finally, we conclude with the implications of the results in terms of the structure of the disk and the respective launching radii of the jets in YSOs. ",
        "conclusions": "We showed as a proof of concept that jet widths derived from numerical simulations extending analytical MHD jet formation models can be very helpful for understanding recently observed jet widths from observations with adaptive optics and space telescopes. However, further aspects have to be investigated in more detail.  An intrinsic feature in all our models with outer truncation is the first bump in the extracted jet width, which complicates the comparison of synthetic and observed jet widths. Only for DG Tau and CW Tau, we could unambiguously find models with outer truncation which fitted the observed jet widths at larger distances. For HN Tau and UZ Tau E, we have no observed jet widths at larger distances, only in the region which is contaminated by the bump. RW Aur shows very small observed jet widths at all scales, which cannot be reproduced by any of our models and runs. The derived truncation radii ($>0.25$ AU) are several times larger than the inner radius of the gaseous disk in T Tauri stars \\citep[0.02-0.07 AU,][]{NCG07}, thus our use of a self-similar disk-wind solution is consistent.  The model with inner truncation gives a synthetic jet width, which is constant and is within the range of observed widths. Another advantage of this model is that its flow velocities of the order of 100 km s$^{-1}$ are closer to observed values. All models with outer truncation have very high velocities ($> 600$ km s$^{-1}$ at our scaling point of $R_{\\rm jet} = 10$ AU and $Z_{\\rm jet} = 100$ AU), several times higher than in that with inner truncation.   Because we originally focused only on the effect of outer truncation, we kept the inner truncation radius constant in our models SC3 and SC5. In order to further explore the ability of inner truncation to reproduce the observations, we have to vary the inner truncation radius in a parameter study. Naively, one would expect that the derived jet widths are not decreasing for decreasing truncation radii as for the outer truncation, but are {\\em increasing}. This, however, has to be tested with new simulations.  In our study, we assumed an inclination of 90$^\\circ$ of the jet, thus projection effects may slightly change our results. We will present the results of such a study in a forthcoming paper.  \\citet{GCF01} showed in their Fig. 1 that the measured jet widths are mainly characterized by the ejection index $\\xi$, defined by \\citet{Fer97}. This is related to the model parameter $x$ of the solution of V00, $\\xi = 2\\,(x - 3/4)$, and because in our simulations $x = 0.75$, we get $\\xi = 0$, which is intrinsic for a standard disk. In the solution of \\citet{Fer97}, $\\xi$ also controls the opening of the field lines, because it is connected to the lever arm $\\lambda$ by $\\lambda = 1 + 1 / (2\\,\\xi)$. \\citet{GCF01} favored cold solutions with values of $\\lambda$ between 50--70. If heating along streamlines is allowed, the relation is broken and also warm solutions of e.g. \\citet{CaF00b} with smaller $\\lambda$ values of 8 and opening of streamlines (the ratio of the maximum radius to the initial launch radius) of about 30 can reproduce the observation. In the solution of V00, we have $\\lambda = G ( \\pi / 2 )^{-2} \\approx 39$ and the maximum opening $G_{\\rm max} / G ( \\pi / 2 ) \\approx 893$. The use of other solutions will therefore highly influence our results in terms of the outer truncation radius.  \\appendix"
    },
    "1005/1005.2981_arXiv.txt": {
        "abstract": "{ {The origin and possible universality of the stellar initial mass function (IMF) is a major issue in astrophysics. One of the main objectives of the $Herschel$ Gould belt survey is to clarify the link between the prestellar core mass function (CMF) and the IMF. We present and discuss the core mass function derived from $Herschel$ data for the large population of prestellar cores discovered with SPIRE and PACS in the Aquila Rift cloud complex at $d \\sim260$~pc. We detect a total of 541 starless cores in the entire $\\sim$11~deg$^2$ area of the field imaged at 70--500~$\\mu$m with SPIRE/PACS. Most of these cores appear to be gravitationally bound, and thus prestellar in nature. Our $Herschel$ results confirm that the shape of the prestellar CMF resembles the stellar IMF, with much higher quality statistics than earlier submillimeter continuum ground-based surveys.} ",
        "introduction": "The question of the origin and possible universality of the IMF, which is crucial for both star formation and galactic evolution, remains a major open problem in astrophysics. The $Herschel$ Space Observatory (Pilbratt et al. 2010), equipped with its two imaging cameras SPIRE (Griffin et al. 2010) and PACS (Poglitsch et al. 2010), provides a unique tool to address this fundamental issue. This problem and other key questions about the earliest phases of star formation and evolution are central to the scientific motivation for the {\\it Herschel} Gould belt survey, which will image the nearby ($d \\leq 0.5$~kpc) molecular cloud complexes of the Gould belt using SPIRE at 250--500~$\\mu$m and PACS at 70--160~$\\mu$m (cf. Andr\\'e \\& Saraceno 2005 and Andr\\'e et al. 2010).  Starting with the 1.2~mm continuum study of the Ophiuchus main cloud by Motte et al. (1998), several ground-based (sub)-millimeter dust continuum surveys of nearby, compact cluster-forming clouds such as $\\rho$~Ophiuchi, Serpens, and Orion~B have uncovered 'complete' (but small) samples of prestellar cores whose associated core mass functions (CMF) resemble the stellar IMF (e.g., Motte et al. 1998; Johnstone et al. 2000; Enoch et al. 2006; Nutter \\& Ward-Thompson 2007; -- see also Alves et al. 2007 and Andr\\' e et al. 2007). Albeit limited by small-number statistics at both ends of the CMF, these findings favor theoretical scenarios according to which the bulk of {\\it the IMF of solar-type stars is largely determined by pre-collapse cloud fragmentation}, prior to the protostellar accretion phase (e.g., Hennebelle \\& Chabrier 2008). The problem of the origin of the IMF may thus largely reduce to a good understanding of the processes responsible for the formation and evolution of prestellar cores within molecular clouds. Apart from limited statistics, another limitation of ground-based submillimeter continuum determinations of the CMF is that they rely on uncertain assumptions about the dust properties (temperature and emissivity). $Herschel$ is now making it possible to dramatically improve the quality of the statistics and to reduce the core mass uncertainties by performing direct measurements of the dust temperatures.  In this Letter, we use the initial results of the Gould belt survey, obtained toward the Aquila Rift complex as part of the science demonstration phase (SDP) of {\\it Herschel}, to present the first CMF derived from $Herschel$ data. We discuss the robustness of the resulting CMF based on several tests and simulations.  The Aquila molecular cloud complex is located at the high Galactic-longitude end of the Aquila Rift, in the neighbourhood of the Serpens star-forming region. Here we adopt a distance of 260~pc for the Aquila star-forming complex. For a detailed introduction of the region and a discussion of its distance, we refer the reader to Bontemps et al. (2010).     \\begin{figure*} \\begin{center} \\begin{minipage}{1.0\\linewidth} \\resizebox{0.54\\hsize}{!}{\\includegraphics[angle=270]{./Fig1a_4_00_EpsPngJpgEps.eps}} \\hspace{2mm} \\resizebox{0.44\\hsize}{!}{\\includegraphics[angle=270]{./Fig1b_4_00_EpsPngJpgEps.eps}} \\end{minipage} \\end{center} \\caption{{\\bf(a)} Column density map derived from SPIRE/PACS observations of Aquila. The subregion referred to as the main subfield in the text is marked by the white rectangle. The cyan triangles mark the positions of the 541 starless cores identified in the entire field. The locations of the HII regions W40 and MWC297/Sh62 are shown. The PDR region around W40 is framed by white polygon, while the dashed square outlines the small region shown in more detail in online Fig.~\\ref{Fig_zooms}a. {\\bf(b)} Same as (a) for the main subfield, with a total of 452 starless cores, marked by cyan triangles.} \\label{Fig_coldens}% \\end{figure*} ",
        "conclusions": "\\subsection{Prestellar nature of the Aquila starless cores}  In this paper, we follow the naming convention that a dense core is called {\\it prestellar} if it is starless {\\it and} gravitationally bound (cf. Andr\\'e et al. 2000, Di Francesco et al. 2007). In other words, prestellar cores represent the subset of starless cores that are most likely to form (proto)stars in the future.  Strictly speaking, spectroscopic observations would be required to derive virial masses for the cores and determine whether they are gravitationally bound or not. However, millimeter line observations in dense gas tracers such as N$_2$H$^+$ show that thermal motions generally dominate over non-thermal motions in starless cores (e.g., Andr\\'e et al. 2007). Assuming that this is indeed the case for the Aquila cores observed here, we may use the critical Bonnor-Ebert (BE) mass, $M_{{\\rm BE}}^{{\\rm crit}}~\\approx~2.4~R_{{\\rm BE}}~a^2 / G$,  as a surrogate for the virial mass, where $R_{{\\rm BE}}$ is the BE radius, $a$ is the isothermal sound speed, and $G$ is the gravitational constant. The critical BE mass may also be expressed as $M_{{\\rm BE}}^{{\\rm crit}} \\approx 1.18 {a^4 \\over G^{3/2}} P_{{\\rm ext}}^{-1/2} $, where $P_{{\\rm ext}}$ is the external pressure, which may be estimated as a function of the column density of the local background cloud, $\\Sigma_{cl}$, as $P_{{\\rm ext}} \\approx 0.88~G~\\Sigma_{cl}^2$ (McKee \\& Tan 2003). For each object, we derived two estimates of the BE mass: (1) $M_{BE}(R_{obs})$ as a function of the observed core radius assuming a gas temperature of 10~K, and (2) $M_{BE}(\\Sigma_{cl})$ as a function of the local background column density measured from both our source-subtracted $Herschel$ images and the near-IR extinction map of Bontemps et al. (2010 -- see also Fig.~\\ref{Fig_coldens_compare}b). We then calculated BE mass ratios of $\\alpha_{\\rm BE} \\equiv \\rm{max}[M_{BE}(R_{obs}), M_{BE}(\\Sigma_{cl})] / M_{obs} $ and selected candidate prestellar cores to be the subset of starless cores for which $\\alpha_{\\rm BE}  \\lesssim 2$. Based on this criterion, 314 (or $\\sim69\\% $) of the 452 starless cores of the main subfield and 341 (or $\\sim63\\% $)  of the 541 starless cores of the entire field were found to be bound and classified as good candidate prestellar cores. These high fractions of bound objects are consistent with the locations of the Aquila starless cores in a mass versus size diagram (see online Fig.~\\ref{Fig_mass_size} for the main subfield and Fig.~4 of Andr\\'e et al. 2010 for the entire field).  The self-gravitating character of most $Herschel$ cores in Aquila is supported further by an independent analysis of their internal column density contrasts. We here define the internal column density contrast of a core as $\\Sigma_{peak}/$$\\bar{\\Sigma}_{core}$, where $\\Sigma_{peak}$ and $\\bar{\\Sigma}_{core}$ are the peak and mean column densities of the core, respectively. Assuming optically thin dust emission at $Herschel$ wavelengths and neglecting any dust temperature/opacity gradient, the internal column density contrast can be estimated from the core intensity values in the same form, as $I_{\\nu}^{\\rm peak}/\\bar{I}_{\\nu}$. The internal column density contrast is closely related to the degree of central concentration of a core defined by Johnstone et al. (2000) as $C =1-(\\bar{\\Sigma}_{core}/\\Sigma_{peak})$. Based on their radial intensity profiles (cf. online Fig.~\\ref{Fig_zooms}b), the Aquila starless cores have a median internal column density contrast $\\sim4$. For comparison, the internal column density contrast is larger than 3.6 for supercritical self-gravitating BE spheres, while it is only 1.5 for a uniform-density, non-self-gravitating object (Johnstone et al. 2000).  As a final check, we also estimated the typical column density contrast of the $Herschel$ cores over the local background. For a critically self-gravitating BE core, the mean column density $\\bar{\\Sigma}_{\\rm BE}^{{\\rm crit}}~\\approx~1.56~(P_{{\\rm ext}}~/ G)^{1/2}$ is expected to exceed the column density of the local background cloud $\\Sigma_{cl}$ by a factor of 1.46, if $P_{{\\rm ext}}~\\approx~0.88~G~\\Sigma_{cl}^2$ (McKee \\& Tan 2003). The candidate prestellar cores identified with $Herschel$ have a median column density contrast of  $\\sim1.5$ over the local background, which is consistent with the conclusion that they are self-gravitating.   \\subsection{Prestellar core mass function in Aquila}  Figure~\\ref{Fig_mass_distr}a shows the mass distribution (CMF) of the 452 starless cores identified in the Aquila main subfield. A  power-law fit to the high-mass end of this CMF gives d$N$/dlog$M$ $\\propto$ $M^{-1.5 \\pm 0.2}$ for $M_{\\rm core} > 2~M_\\odot$, which is very close to the Salpeter power-law IMF (d$N/$dlog$M$ $\\propto$ $M^{-1.35}$ -- Salpeter 1955). The CMF of the 541 starless cores identified in the entire field, shown by Andr\\'e et al. (2010), has a very similar shape.  The robustness of the derived CMF was tested by selecting several subsets of sources, based e.g., on their physical nature or locations in the SPIRE/PACS maps. In particular, we compared the mass spectrum of starless and prestellar cores. Figure~\\ref{Fig_mass_distr}a was then compared with Fig.~\\ref{Fig_mass_distr}b, which shows the differential mass function of 314 bound cores in the main subfield. We obtained the same best-fit power law fit to the high-mass end (d$N$/dlog$M$ $\\propto$ $M^{-1.45 \\pm 0.2}$) as for Fig.~\\ref{Fig_mass_distr}a. Only the low-mass end changed when selecting the subset of candidate prestellar cores.  We also selected subsets of starless cores based on their locations in the Aquila field. Along lines of sight to the HII region W40, the associated photon-dominated region (PDR, Shuping et al. 1999) is the source of very strong extended background emission at all infrared wavelengths. The strong background emission in the {\\it Spitzer} 24~$\\mu$m and PACS 70~$\\mu$m images makes it more difficult to discriminate between YSOs and compact starless cores, implying that our census of starless cores is less reliable in this area. We thus constructed another version of the Aquila main subfield CMF, excluding the 83 cores identified toward the PDR region. The excluded region of unusually high infrared background emission (see in Fig.~\\ref{Fig_coldens}a) was defined using the dust temperature map shown in Bontemps et al. (2010). The high-mass end of this CMF can be fitted by a very similar power law to that of Fig.~\\ref{Fig_mass_distr}a: d$N$/dlog$M$ $\\propto$ $M^{-1.5 \\pm 0.3}$.  Thanks to the large number of starless cores identified with $Herschel$ in Aquila (541 cores in the entire field), we have been able to consider several core subsamples and construct a statistically meaningful CMF in each case. We confirm that the shape of the prestellar CMF resembles the stellar IMF, using data of far higher quality statistics than earlier submillimeter ground-based surveys and more accurate core masses. Based on simulations, we conclude that our mass distributions are robust and do not depend strongly on different sets of extracted sources.  The column density maps shown in Fig.~\\ref{Fig_coldens} and online Fig.~5 illustrate the tight correlation between the spatial distribution of the prestellar cores and the filamentary structure revealed by the $Herschel$ images (Men'shchikov et al. 2010). The difference between our two SDP fields (Aquila Rift and Polaris Flare), in terms of filamentary structure and core mass distribution, is highlighted and discussed in Andr\\'e et al. (2010)."
    },
    "1005/1005.0984_arXiv.txt": {
        "abstract": "{Stellar rotation is a crucial parameter driving stellar magnetism, activity and mixing of chemical elements. Measuring rotational velocities of young stars can give additional insight in the initial conditions of the star formation process. Furthermore, the evolution of stellar rotation is coupled to the evolution of circumstellar disks. Disk-braking mechanisms are believed to be responsible for rotational deceleration during the accretion phase, and rotational spin-up during the contraction phase after decoupling from the disk for fast rotators arriving at the ZAMS. On the ZAMS, stars get rotationally braked by solar-type winds.} {We investigate the projected rotational velocities $v\\sin i$ of a sample of young stars with respect to the stellar mass and disk evolutionary state to search for possible indications of disk-braking mechanisms. Furthermore, we search for signs of rotational spin-up of stars that have already decoupled from their circumstellar disks.} {We analyse the stellar spectra of 220 nearby (mostly $<100$\\,pc) young (2--600\\,Myr) stars for their $v\\sin i$, stellar age, H$\\alpha$ emission, and accretion rates. The stars have been observed with FEROS at the 2.2m MPG/ESO telescope and HARPS at the 3.6m telescope in La Silla, Chile. The spectra have been cross-correlated with appropriate theoretical templates. We build a new calibration to be able to derive $v\\sin i$ values from the cross-correlated spectra. Stellar ages are estimated from the \\ion{Li}{i} equivalent width at 6708\\,\\AA. The equivalent width and width at 10\\% height of the H$\\alpha$ emission are measured to identify accretors and used to estimate accretion rates $\\dot{M}_\\mathrm{acc}$. The $v\\sin i$ is then analysed with respect to the evolutionary state of the circumstellar disks to search for indications of disk-braking mechanisms in accretors.} {We find that the broad $v\\sin i$ distribution of our targets extends to rotation velocities of up to more than 100\\,kms$^{-1}$ and peaks at a value of $7.8\\pm1.2$\\,kms$^{-1}$, and that $\\sim70\\%$ of our stars show $v\\sin i<30$\\,kms$^{-1}$. Furthermore, we can find indications for disk-braking in accretors and rotational spin-up of stars which are decoupled from their disks. In addition, we show that a number of young stars are suitable for precise radial-velocity measurements for planet-search surveys.} {} ",
        "introduction": "\\label{sec_intro}  Rotation is one of the most important kinematic properties of stars, % giving rise to stellar magnetism and mixing of chemical elements. Stars form from rotating molecular cloud cores, preserving only a very minor fraction of their initial angular momentum (e.g., Palla 2002, Lamm et al. 2005). The initial angular momentum of rotating molecular cloud cores is about 4--5 orders of magnitude higher than that of the stars that eventually form in this cloud core (e.g., Bodenheimer 1989). Stellar formation models must account for that and several mechanisms are discussed, e.g., magnetic braking or magnetocentrifugally driven outflows. However, there is still no consens as on whether there is one dominant process for dispersing angular momentum during the entire star formation process or which process dominates at what evolutionary stage after the formation of a star-disk system (e.g., Palla 2002).  The mechanisms believed to be responsible for efficient loss of stellar angular momentum after the formation of a star-disk system involve transfer of angular momentum along magnetic field lines that connect the stellar surface with the disk, either onto the disk (so-called 'disk-locking', e.g. Camenzind 1990, K\\\"onigl 1991) or into stellar winds originating at that boundary (Matt \\& Pudritz 2005 and 2008). Both models predict that the stellar magnetic field threads the circumstellar disk, accretion of disk material onto the star occurs along the field lines, and magnetic torques transfer angular momentum away from the star. In the disk-locking model, the angular velocity of the star is then locked to the Keplerian velocity at the disk boundary. For reasonable magnetic field strengths and accretion rates $\\sim$\\,10$^{-9}$\\,M$_{\\odot}$/yr, both models account for the observed rotational periods of classical T~Tauri stars (e.g., Armitage \\& Clarke 1996). Both models can be observationally distinguished only by the absence or presence of stellar winds originating from the boundary. When the accretion process ends and there is no longer an efficient way to disperse angular momentum, the star can spin-up during the contraction phase (see, e.g., Lamm et al. 2005, Bouwman et al. 2006). Once arrived on the ZAMS, stars undergo additional rotational braking by magnetic winds, irrespective of whether they arrived as slow or fast rotators (e.g., Skumanich 1972, Palla 2002, Wolff et al. 2004).  If this picture holds true, very young stars are expected to rotate slower than slightly older stars that have already decoupled from their disks. Three different classes of young stars are expected to be distinguishable, which can be interpreted as a kind of evolutionary sequence: (i) slow rotating stars which still accrete, (ii) slow to intermediate rotating stars which do not longer accrete and are about to gradually spin-up due to contraction and (iii) fast rotators without disks. However, one has to keep in mind, that all star forming regions show an enormous spread in rotation periods and that differences in rotation velocities between different star forming regions can also be caused by the initial conditions of star formation and the braking time-scale of the star-disk system (e.g., Stassun et al. 1999, Lamm et al. 2005, Herbst et al. 2007, Nguyen et al. 2009).  Several surveys of stellar rotation periods and projected rotational velocities of very young stars indeed found relatively slow rotators (e.g., Herbst et al. 2002; Rebull et al. 2004, 2006; Sicilia-Aguilar et al. 2005), which points to effective braking mechanisms, whereas other surveys did not find evidence for rotational braking of very young stars  (e.g., Makidon et al. 2004; Nguyen et al. 2009). Nevertheless, rotational velocity measurements of stars with circumstellar disks and ongoing accretion in associations such as $\\eta$\\,Cha, TW~Hydrae, and NGC\\,2264 support a disk-locking mechanism for the removal of angular momentum from the star (Lamm et al. 2005; Bouwman et al. 2006; Jayawardhana et al. 2006; Fallscheer \\& Herbst 2006). This rotational braking due to disk-locking is mainly visible in very young clusters with an age $<$\\,3\\,Myr. These stars are expected to spin-up by a factor of $\\approx$\\,3 due to contraction after being magnetically disconnected from the circumstellar disk. In fact, a large fraction of PMS stars have been observed that evolve at nearly constant angular velocity through the first 4\\,Myr (Rebull et al. 2004). In clusters at the age of $\\sim$\\,10\\,Myr, signs of stellar velocity spin-up can be seen (Rebull et al. 2006).  For our radial-velocity survey of very young stars, SERAM (\\emph{Search for Exoplanets with Radial-velocity At the MPIA}, Setiawan et al. 2008), we started to measure the $v\\sin i$ of stars with ages of 2--600\\,Myr and selected slow rotators for high-precision radial velocity measurements to search for planets. These $v\\sin i$ measurements are a necessary prerequisite for such an RV survey, because the projected rotational velocity of a star limits the accuracy of the radial velocity measurement due to broadening of the spectral lines. Furthermore, rotational velocities play a crucial role in describing and understanding chromospheric and coronal activity, which is related to radial velocity jitter due to stellar spots and scales with rotational velocities (e.g., Pallavicini et al. 1981, Noyes et al. 1984, Saar \\& Donahue 1997). Thus, measuring rotational velocities and periods of young stars is very important to understand stellar formation, evolution and activity.  This paper is structured as follows. In section \\ref{obs}, we describe the target selection, observations, and data reduction. The parameter estimation is described in Section~\\ref{sec_dataana} and the calibration of $v\\sin i$ is described in Appendix~\\ref{sec_vm}. In Section~\\ref{res}, we present the results of our analysis and search for indications of disk-braking and rotational spin-up. Finally, we summarize our results in Section \\ref{con}. ",
        "conclusions": "\\label{con}  We analysed 229 young and nearby stars for their projected rotational velocity $v\\sin i$, stellar age, and accretion signatures. The stars were part of our initial RV planet search target list in the SERAM project and have been observed with FEROS at the 2.2m MPG/ESO telescope and HARPS at the 3.6m telescope in La Silla, Chile. For stars showing broad H$\\alpha$ emission, we checked other accretion indicators, like the EW of \\ion{He}{i} at 5876\\,\\AA, to identify those stars that still accrete material from a circumstellar disk. For these stars, the veiling has also been measured. The age of the stars has been derived from the EW of \\ion{Li}{i} at 6708\\,\\AA{}.  We calculated $v\\sin i$ of our 220 target stars and 9 spectroscopic companions by measuring the width of the CCF of the stellar spectra. To do this, we cross-correlated the stellar spectra with theoretical templates of similar spectral type. The CCF has been fitted with a Gaussian and the conversion to $v\\sin i$ was done with our new calibration (Equation~\\ref{eq_sigma0}) for both FEROS and HARPS spectra.  The main conclusions from this work can be summarized as follows: \\begin{itemize} \\item{The mean stellar age of our sample is 30\\,Myr with a spread (width of Gaussian fit) of 20\\,Myr. 43 stars are younger than 10\\,Myr.} \\item{The distribution of $v\\sin i$ for all target stars peaks at 7.8$\\pm$1.2\\,kms$^{-1}$, with the vast majority having projected rotational velocities $<30$\\,kms$^{-1}$.} \\item{We found indication for rotational braking due to disk-locking, because the accreting stars in our sample rotate significantly slower ($\\langle v\\sin i\\rangle =10\\pm1$\\,kms$^{-1}$) than the non-accreting stars with more evolved disks ($\\langle v\\sin i\\rangle =15\\pm2$\\,kms$^{-1}$). The only 2 fast rotating ($v\\sin i \\sim 38$\\,kms$^{-1}$) accretors in our sample are DI\\,Cha and CR\\,Cha. This might point to differences in time spent for disk-braking or to different initial conditions in different star formation regions, but the significance of this statement is hampered by low number statistics.} \\item{This $v\\sin i$ evolution also provides indication that stars undergo rotational spin-up after being decoupled from their circumstellar disk and evolving along the Pre-Main-Sequence track. For more evolved stars with debris or no disks and an age\\,$>30$\\,Myr, $\\langle v\\sin i\\rangle$ decreases again to 10\\,$\\pm$\\,1\\,kms$^{-1}$. We conclude that these stars have been rotationally braked on the ZAMS, e.g., by magnetic winds. } \\item{We estimated the maximum rotational period of our target stars. The period distribution of the whole sample peaks at 2.5\\,$\\pm$\\,0.3\\,days, a similar result as found for young stars by Cieza \\& Biliber (2007).} \\item{Due to disk-braking, young and accreting stars can rotate sufficiently slow to serve as suitable targets for RV planet searches, since the accuracy of radial velocity measurements depends on the stellar $v\\sin i$. When accretion ends after $\\sim$5--10\\,Myr and the stars decouple from their disk, they tend to spin up as they contract, such that they are usually no longer suitable for precise RV measurements. This changes only when the stars arrive on the ZAMS and are slowed down again by wind braking at ages of $\\sim$30\\,Myr.} \\end{itemize}"
    },
    "1005/1005.3569_arXiv.txt": {
        "abstract": "\\noindent On 1998 November 14, Saturn and its rings occulted the star GSC 0622-00345.  We observed atmospheric immersion with NSFCAM at the National Aeronautics and Space Administration's Infrared Telescope Facility on Mauna Kea, Hawaii.  Immersion occurred at 55{\\decdegree}5 S planetocentric latitude.  A 2.3 {\\micron}, CH\\sb{4}-band filter suppressed reflected sunlight.  Atmospheric emersion and ring data were not successfully obtained.  We describe our observation, light-curve production, and timing techniques, including improvements in aperture positioning, removal of telluric scintillation effects, and timing.  Many of these techniques are known within the occultation community, but have not been described in the reviewed literature.  We present a light curve whose signal-to-noise ratio per scale height is 267, among the best ground-based signals yet achieved, despite a disadvantage of up to 8 mag in the stellar flux compared to prior work.  \\if\\submitms y \\else \\comment{\\hfill\\herenote{DRAFT of {\\today} \\now}.} \\fi ",
        "introduction": "\\label{intro}  Stellar occultations by planets offer among the highest spatial resolution of any astronomical observing technique, competing even with atmospheric descent probes for certain measurements.  From Earth, spatial resolution is limited by the projected stellar beam size and Fresnel diffraction, the latter giving \\sim1 km resolution for solar system giant planets for visible and near-IR wavelengths.  This is sufficient to reveal fine structure in a planetary ring or propagating gravity (buoyancy) waves in an atmosphere.  However, the occultation technique is demanding, and failing to optimize even a single parameter of an observation can significantly reduce the quality of the derived light curve.  Occultation observers have thus developed numerous specialized techniques to improve their data.  Many of these have never been described in the literature.  On 1998 November 14, Saturn and its rings occulted GSC 0622-00345, as predicted by \\citet{BoshMcDonald1992ajsatocs}.  We obtained a light curve for atmospheric immersion, based on infrared imaging observations at the NASA Infrared Telescope Facility (IRTF) on Mauna Kea, HI.  Figure \\ref{geometry} shows our viewing circumstances and Table \\ref{occtab} gives parameters of the event.  Occultation cadences are typically from a few to 10 Hz, and in the case of lunar occultations they can be much higher.  These rates are sufficient for the temporal resolution of flux variations caused by scintillation (variable refractive focusing and defocusing of the stellar beam) in Earth's atmosphere.  This is typically the limiting noise source for occultations by bright stars.  During our observations, we monitored Rhea, which experienced approximately the same scintillation effects as the star, just arcseconds away.  We were thus able to use measurements of Rhea's (assumed constant) flux to compensate for scintillation.  This, in combination with other techniques, produced a light curve of very high quality.  Here we present a full accounting of the methods we employed to observe the event and derive the light curve, including those previously undocumented.  Section \\ref{obslcsec} describes observing techniques, data reduction methods, and the light curve.  It also discusses an attempt to apply optimal photometry to light curve production and thoughts on improving instrument interfaces to avoid failures such as our lost egress observation.  Section \\ref{timingsec} presents a method for independently verifying the system timing solution.  Section \\ref{conclsec} offers brief conclusions.  We include the light curve and software for analyzing occultation timing data as electronic supplements.  Paper II \\citep{HarringtonEtal2010apjSatoc98II} uses the derived light curve to investigate the atmosphere of Saturn. ",
        "conclusions": "\\label{conclsec}  We have presented a light curve based on IRTF observations of the 1998 November 14 occultation of GSC 0622-00345 by Saturn.  The light curve has a per-frame S/N of 66 and a per-scale-height S/N of 267, placing it among the best ground-based atmospheric light curves.  We have described the observational and analytic methods used to derive the light curve, in some cases for the first time in the literature, although these techniques are known among experienced occultation observers.  Methods that improved the photometry include using Rhea as a scintillation and pointing standard, spatially subsampling the data by a factor of 10 for accurate aperture placement, and removing a template of scattered light.  Separate, identical, measurements of a flashing LED placed in the telescope beam provided an independent measure of the frame rate, when fit by a detailed model of the chip's sensitivity throughout the imaging cycle.  Optimal photometry shows promise at reducing the photometric noise level as the star dims, but requires a PSF standard that is substantially brighter than the occultation star.  Occultation cameras with wider fields of view (but still able to operate with accurate timing at high frame rates) would facilitate finding bright PSF standards."
    },
    "1005/1005.1877_arXiv.txt": {
        "abstract": "{We present the observations of the starburst galaxy \\object{M82} taken with the \\textit{Herschel} SPIRE Fourier Transform Spectrometer. The spectrum (194--671 $\\mu$m) shows a prominent CO rotational ladder from $J$ = 4--3 to 13--12 emitted by the central region of \\object{M82}. The fundamental properties of the gas are well constrained by the high $J$ lines observed for the first time. Radiative transfer modeling of these high-S/N $^{12}$CO and $^{13}$CO lines strongly indicates a very warm molecular gas component at $\\sim 500$~K and pressure of $\\sim 3 \\times 10^{6}$~K~cm$^{-3}$,  in good agreement with the H$_{2}$ rotational lines measurements from \\textit{Spitzer} and \\textit{ISO}. We suggest that this warm gas is heated by dissipation of turbulence in the interstellar medium (ISM) rather than X-rays or UV flux from the straburst. This paper illustrates the promise of the SPIRE FTS for the study of the ISM of nearby galaxies.} ",
        "introduction": "Starburst galaxies provide us with the opportunity to study star formation and its effect on the interstellar medium (ISM) in extreme environments. These galaxies combine large central gas concentrations and high ionizing radiation fields, resulting in bright molecular, neutral and ionized gas emission lines.  At a distance of 3.9 Mpc \\citep{sakai99}, \\object{M82} is the most well-studied starburst galaxy in the local universe, and it is widely used as a starburst prototype in cosmological studies. Its infrared luminosity \\citep[$5.6\\times 10^{10}$ L$_\\odot$,][]{sanders03} corresponds to a star-formation rate of 9.8 M$_\\odot$ yr$^{-1}$, which has almost certainly been enhanced by its interaction with \\object{M81} and \\object{NGC 3077} \\citep{yun93}. With a reported molecular gas content of $1.3 \\times 10^9$ M$_\\odot$ \\citep{walter02}, its bright emission lines of CO and other molecules allow us to study its ISM in great detail \\citep{shen95,walter02,ward03}.  Far-infrared fine structure lines were used to constrain the physical properties of the ionized gas and photo-dissociation regions (PDRs) in \\object{M82}. \\citet{colbert99} found that the ionized gas emission can be reproduced with a 3--5 Myr old instantaneous starburst and a gas density of 250 cm$^{-3}$, while the PDR component is best fit with a density of 2\\,000 cm$^{-3}$, in pressure equilibrium with the ionized phase.  Stellar evolution and photoionization models \\citep{forster03} indicate a series of a few, Myr-duration starbursts with a peak of activity 10 Myr ago in the central regions, and 5 Myr ago in the circumnuclear ring. Models of the PDR and molecular emission as a set of non-interacting hot bubbles driving spherical shells of swept-up gas into a surrounding uniform medium also predict a starburst age of 5--10 Myr, but fail to match the observed far-infrared luminosity \\citep{yao09}.  The strengths of the CO lines place fundamental constraints on the physical properties of the molecular gas. \\cite{tilanus91} fitted \\coone\\ and \\cotwo\\ lines from the central starburst up to $J$ = 3--2 with a single-component model with temperatures of 30--55 K and densities of 3--7 $\\times 10^3$ cm$^{-3}$. \\citet{wild92} used lines up to the CO $J$ = 6--5 transition to refine these parameters to 40--50 K and $\\sim 10^4$ cm$^{-3}$, while HCN and HCO$^+$ lines suggested densities greater than $3\\times 10^5$ cm$^{-3}$ are present. \\citet{petitpas00} showed evidence for a temperature or density gradient across the starburst region. \\citet{weiss05} showed that CO emission up to $J$ = 3--2 is dominated by more extended regions while higher $J$ transitions originate in the central disk.  In this paper, we present observations of \\object{M82} with \\textit{Herschel} \\citep{pilbratt10} using the SPIRE Fourier Transform Spectrometer (FTS) \\citep{griffin10}, which measures the complete far-infrared spectrum from 194 to 671 $\\mu$m. This spectral region is particularly interesting for probing the peak of the CO spectral line energy distribution (SLED) in gas-rich galaxies. The wealth of lines across a continuous spectral region allows for unprecedented precision in modeling the physical and chemical properties of the molecular ISM. Here, we focus on the measurement and analysis of the CO rotational transitions from the central starburst in \\object{M82}. ",
        "conclusions": "We have presented the \\textit{Herschel}-SPIRE spectroscopic observations of the starburst galaxy \\object{M82}. The spectra show a prominent CO emission-line ladder along with \\ion{C}{I} and \\ion{N}{II} lines. Radiative transfer modeling of  CO lines clearly indicates a warm gas component at $\\sim$500~K in addition to the cold ($\\sim$30~K) component derived from ground-based studies. The properties of the warm gas are strongly constrained by the high $J$ lines, observed here for the first time. The temperature and mass of warm gas agree with the H$_2$ rotational lines observations from \\textit{Spitzer} and \\textit{ISO}. At this temperature H$_{2}$ is the dominant coolant instead of CO, and we argue that turbulence from stellar winds and supernovae may be the dominant heating mechanism."
    },
    "1005/1005.1931_arXiv.txt": {
        "abstract": "According to the no-hair theorem, all astrophysical black holes are fully described by their masses and spins. This theorem can be tested observationally by measuring (at least) three different multipole moments of the spacetimes of black holes. In this paper, we analyze images of black holes within a framework that allows us to calculate observables in the electromagnetic spectrum as a function of the mass, spin, and, independently, the quadrupole moment of a black hole. We show that a deviation of the quadrupole moment from the expected Kerr value leads to images of black holes that are either prolate or oblate depending on the sign and magnitude of the deviation. In addition, there is a ring-like structure around the black-hole shadow with a diameter of $\\sim$10 black-hole masses that is substantially brighter than the image of the underlying accretion flow and that is independent of the astrophysical details of accretion flow models. We show that the shape of this ring depends directly on the mass, spin, and quadrupole moment of the black hole and can be used for an independent measurement of all three parameters. In particular, we demonstrate that this ring is highly circular for a Kerr black hole with a spin $a\\lesssim0.9M$, independent of the observer's inclination, but becomes elliptical and asymmetric if the no-hair theorem is violated. Near-future very-long baseline interferometric observations of Sgr A* will image this ring and may allow for an observational test of the no-hair theorem. ",
        "introduction": "According to the no-hair theorem, the external spacetimes of black holes are uniquely characterized by their masses and spins (Israel 1967, 1968; Carter 1971, 1973; Hawking 1972; Robinson 1975). This theorem requires that the cosmic censorship conjecture (Penrose 1969) holds and that the exterior spacetime is free of closed time-like curves. Under these assumptions, all astrophysical black holes should be fully described by the Kerr metric (Kerr 1963).  It is widely accepted that the universe contains an abundance of black holes as inferred from the observations of the centers of nearby galaxies (e.g., Tremaine et al.\\ 2002), of our own galactic center (Ghez et al. 2008; Gillessen et al. 2009), and of many galactic binaries (e.g., McClintock \\& Remillard 2006). Nonetheless, the factual existence of an event horizon is yet unproven and has only been inferred indirectly (e.g., Narayan, Garcia, \\& McClintock 1997, 2001; Narayan \\& Heyl 2002; McClintock, Narayan, \\& Rybicki 2004; see also Psaltis 2006). Alternative explanations for the nature of these objects have been suggested, which include naked singularities (Manko \\& Novikov 1992), exotic stellar objects (Friedberg, Lee, \\& Pang 1987; Mazur \\& Mottola 2001; Barcel\\'{o} et al. 2008), as well as a breakdown of general relativity itself on horizon scales (e.g., Yunes \\& Pretorius 2009; c.f. Psaltis et al. 2008).  A test of the no-hair theorem can both identify the observed dark compact objects with Kerr black holes and verify the validity of general relativity in the strong-field regime. Indeed, within general relativity, if a compact object is not a Kerr black hole, then its external spacetime will not satisfy the no-hair theorem. Alternatively, if general relativity is not valid in the strong-field regime, the external spacetime of a compact object that is surrounded by a horizon may violate the no-hair theorem (see, however, Psaltis et al.\\ 2008).  Mass and spin are the first two multipole moments of a black-hole spacetime. If the no-hair theorem is correct, then all higher multipole moments only depend on the mass and spin, and any deviation from the Kerr moments has to be zero. Consequently, the no-hair theorem can be tested by measuring (at least) three multipole moments of such a spacetime (Ryan 1995).  In part I of this series of papers (Johannsen \\& Psaltis 2010; hereafter Paper I), we investigated a framework for testing the no-hair theorem with observations of compact objects in the electromagnetic spectrum. Based on a quasi-Kerr spacetime that contains an independent quadrupole moment (Glampedakis \\& Babak 2006), we analyzed in detail the spacetime properties that are critical for such observations as a function of the mass, spin, and quadrupole moment. We showed that already very moderate changes of the quadrupole moment lead to significant alterations of various quantities that determine observables. In particular, we explored the effect of changing the quadrupole moment on the locations of the innermost stable circular orbit (ISCO) and of the circular photon orbit, as well as on the lensing and redshift of photons.  There has been, already, substantial work on potential tests of the no-hair theorem with observations of gravitational waves from extreme mass-ratio inspirals (Ryan 1995, 1997a, 1997b; Barack \\& Cutler 2004, 2007; Collins \\& Hughes 2004; Glampedakis \\& Babak 2006; Gair et al.\\ 2008; Li \\& Lovelace 2008; Apostolatos et al.\\ 2009; Vigeland \\& Hughes 2010).  In this series of papers, we show that observations of black holes in the electromagnetic spectrum may also allow for a clean test of the no-hair theorem. In particular, we identify different observables that probe the quadrupole moments of the spacetimes but depend very weakly on the usual astrophysical complications, such as the flow geometry, the mode of emission, and the variability of the accretion flows that generate the photons we detect from black holes.  Imaging observations of accreting black holes at (sub$-$)mm wavelengths using very-long baseline interferometry (VLBI) promise to enable unprecedented views of the vicinities of black-hole horizons. Recent VLBI-observations along only three baselines resolved Sgr~A$^*$, the black hole in the center of the Milky Way, on a scale comparable to its event horizon and provided evidence for sub-horizon scale structures (Doeleman et al.\\ 2008) as well as for the presence of an event horizon (Broderick, Loeb, \\& Narayan 2009). Far greater resolution can be achieved by adding either existing or planned telescopes located at various places on the Earth (Fish \\& Doeleman 2009). The black hole in the center of M87 (Broderick \\& Loeb 2009) as well as a small number of other nearby supermassive black holes (Psaltis 2008) offer additional targets for horizon-scale imaging that can be utilized in the near future.  Imaging observations are expected to be able to resolve the shadows of black holes and lead to the determination of their spins and inclinations (Falcke et al. 2000; Broderick \\& Loeb 2005, 2006; Fish \\& Doeleman 2009). Templates for images of accretion flows around Kerr black holes within general relativity that are suitable to these observations have been reported by a number of authors (Bardeen 1973; Speith et al. 1995; Fanton et al. 1997; Falcke et al. 2000; Takahashi 2004; Beckwith \\& Done 2004, 2005; Dexter \\& Agol 2009; Broderick \\& Loeb 2005, 2006; Yuan et al. 2009). Recently, Bambi \\& Freese (2009) explored the possibility of using black-hole images to test whether black holes violate the Kerr bound $a\\le M$.  In this paper, we study the properties of the images of compact objects that violate the no-hair theorem using the quasi-Kerr formalism we developed in Paper I.  We calculate numerically the mapping between locations in the vicinity of a black hole and positions in the observer's sky using the mass, spin, and quadrupole moment of the spacetime as independent parameters. We investigate the impact of varying the quadrupole moment on the properties of this mapping and show that the images of the accretion flows around compact objects that violate the no-hair theorem are expected to have prolate or oblate geometries.  Measuring the spacetime moments from the images of an accretion flow will be, of course, very model dependent and limited by our lack of understanding of the intrinsic geometry of the flow itself. For example, prolate images of the inner accretion may be the result of resolving the formation region of a jet and not of a violation of the no-hair theorem (see, e.g., Broderick \\& Loeb 2009). Moreover, a measurement of the spin from an image of the shadow alone is difficult (e.g., Falcke et al. 2000; Takahashi 2004) and might require complementary observations such as a multiwavelength study of polarization (Broderick \\& Loeb 2006; see also Schnittman \\& Krolik 2009, 2010).  Additionally, accretion flows are very turbulent and variable at timescales much shorter than the rotation period of the Earth, which sets the characteristic integration time for an interferometric imaging observation. If this time variability is produced by a highly coherent, orbiting inhomogeneity in the accretion flow, it may allow measuring the properties of the compact object via non-imaging techniques (Doeleman et al.\\ 2009). The variability, however, will limit and may prohibit altogether the ability of obtaining a clean image of the accretion flow.  The images of optically thin accretion flows around black holes, however, reveal a characteristic bright ring at the projected radius of the circular photon orbit along null geodesics (Beckwith \\& Done 2005) with properties that remain constant even as the underlying accretion flows are highly variable. This bright ring is the result of the light rays that orbit around the black hole many times before they reach the distant observer, and, therefore, have a much larger path length through the optically thin accretion flow. These photons can make a significant contribution to the total disk emission and produce higher-order images (Cunningham 1976; Laor, Netzer, \\& Piran 1990; Viergutz 1993; Bao, Hadrava, \\& {\\O}stgaard 1994; ${\\rm \\check{C}ade\\check{z}}$, Fanton, \\& Calvani 1998; Agol \\& Krolik 2000; Beckwith \\& Done 2005).  We use our formalism to show that the bright emission ring is circular for a Schwarzschild black hole and remains nearly circular for Kerr black holes. On the other hand, if the quadrupole moment is left as an independent parameter, the ring shape changes significantly and becomes asymmetric. The degree of asymmetry is a direct measure of the violation of the no-hair theorem. We show that the diameter of the ring depends only very weakly on the spin and quadrupole moment of the black hole and can be used to directly measure the mass of the object. In addition, the ring is displaced off center in the image plane in the case of rotating black holes (Beckwith \\& Done 2005; see, also, Takahashi 2004), and we show that the displacement is a direct measure of the object's spin, modulo the disk inclination.  In Section~2 we briefly review the framework for testing the no-hair theorem with observations in the electromagnetic spectrum. We simulate images of black holes and of the photon rings in Sections~3 and 4, respectively, and show how they depend on the mass, spin, and quadrupole moment of a given black hole. In Section~5 we quantify these dependencies before we discuss our results and their implications for future observations in Section~6. ",
        "conclusions": "In Paper~I, we proposed a new framework for testing the no-hair theorem with observations of black holes in the electromagnetic spectrum. We formulated our tests based on a quasi-Kerr metric (Glampedakis \\& Babak 2006), which deviates smoothly from the Kerr metric in the quadrupole moment. Since the no-hair theorem admits exactly two independent multipole moments for a black hole, a measurement of these three moments will allow us to test the no-hair theorem.  In this paper, we calculated numerically the mapping between different locations in the accretion flow around a quasi-Kerr black hole and positions in the image plane of a distant observer. Our calculations allowed us to study the potential of using imaging observations of black holes, that will become available in the near future, in order to test the no-hair theorem.  We argued that the expected image of an accretion flow will be characterized by a bright emission ring generated by light rays that circle multiple times around the event horizon before emerging towards the observer. We identified the ring diameter as a direct measure of the mass of the black hole, and we quantified the dependence of the displacement and the asymmetry of the ring on the spin and the quadrupolar parameter as well as on the disk inclination. For a given inclination angle, a measurement of the displacement and the asymmetry directly measures the spin and the quadrupolar parameter of the system, respectively. The asymmetry itself provides a direct measure of the violation of the no-hair theorem.  It is important to emphasize here that only the relative displacement and asymmetry of the ring (i.e., measured in units of the ring diameter) and not their absolute values are necessary in inferring the spin and quadrupole moment of the black-hole spacetime. As a result, the outcome of such an observation does not depend on the distance to the black hole, which might not be known accurately. On the other hand, the angular diameter of the photon ring is proportional to the mass of the black hole. This can lead to an accurate measurement of the mass of the black hole if the distance is known, or else of the distance to the black hole, if its mass is known from, e.g., dynamical observations.  Sgr A*, the black hole in the center of the Milky Way, is the ideal candidate for a test of the no-hair theorem due to its high brightness, large angular size, and relatively unimpeded observational accessibility. Recent VLBI observations (Doeleman et al. 2008) resolved Sgr A* on horizon scales. Incorporating additional baselines to the VLBI network will lead to the first images of Sgr~A$^{*}$ within the next few years (Fish \\& Doeleman 2009). The emission from Sgr~A$^{*}$ at sub-mm wavelengths is optically thin, the size of the scattering ellipse is $\\lesssim1M$, and the resolution of a VLBI image at this wavelength is also comparable to $\\simeq1M$ (Doeleman et al.\\ 2008; Fish \\& Doeleman 2009). The smearing of the images due to scattering and the finite resolution of the array will, therefore, not preclude measuring the position and asymmetry of the photon ring to an accuracy that is adequate in providing a quantitative test of the no-hair theorem.  Observations of the orbits of stars in the vicinity of the black hole have provided an independent measurement of its mass (Ghez et al.\\ 2008; Gillessen et al.\\ 2009). Perhaps more importantly, the same observations may lead in the future to an independent measurement of the spin and orientation of the black hole, as well as to a complementary test of the no-hair theorem (Will 2008; Merritt et al.\\ 2010)."
    },
    "1005/1005.3105_arXiv.txt": {
        "abstract": "Roche-lobe overflow and common envelope evolution are very important in binary evolution, which is believed to be the main evolutionary channel to hot subdwarf stars. The details of these processes are difficult to model, but adiabatic expansion provides an excellent approximation to the structure of a donor star undergoing dynamical timescale mass transfer. We can use this model to study the responses of stars of various masses and evolutionary stages as potential donor stars, with the urgent goal of obtaining more accurate stability criteria for dynamical mass transfer in binary population synthesis studies. As examples, we describe here several models with the initial masses equal to $1~M_\\odot$ and $10~M_\\odot$, and identify potential limitations to the use of our results for giant branch stars. ",
        "introduction": "\\label{sec1}  The Fourth Meeting on Hot Subdwarf Stars and Related Objects was convened in Shanghai, China from 19th to 24th, July. Recent discoveries and developments in both theory and observation of hot subdwarfs and related objects were reported, and many unsolved problems were discussed. As we all know, hot subdwarf stars are extreme horizontal-branch stars or related objects. They may dominate the UV-upturn of early-type galaxies and they exist in both the field of our Galaxy and its globular cluster system. About half of subdwarf B (sdB) stars are binaries \\citep{max01,saf01}; binary evolution is obviously important in their formation, as the observed systems are too compact to have avoided past mass transfer. Binary population synthesis models \\citep{han02,han03,han08} explain naturally the sdB binary fractions and the UV-upturn of early-type galaxies via: (a) one or two phases of common envelope (CE) evolution, (b) stable Roche-lobe overflow (RLOF), and (c) the merger of two He-core white dwarf stars (WDs).  RLOF and CE evolution are also very important in the formation of other binary systems, e.g., cataclysmic variables, X-ray binaries, double white dwarfs and binary neutron stars. Those systems containing compact objects are among the most energetic and rapidly variable sources known. Unfortunately, we know very little about the details of CE evolution, but dynamically unstable RLOF appears to be the trigger that launches it. The threshold conditions for dynamical mass transfer now in common use in binary population synthesis calculations are based on polytropic models for rapid mass loss process \\citep[and references therein]{hjel87}. These polytropic studies provide useful qualitative insights into RLOF and CE evolution, but they omit much relevant physics, and fail to address many advanced evolutionary stages of interest. Detailed studies of binary evolution \\citep{han02,chen08} reveal a need for more realistic determinations of the threshold conditions for dynamical mass transfer.  With the motivation and cautions above, we set out to study stellar rapid mass loss based on the pioneering work of \\citet{hjel89a,hjel89b}. We describe the basic assumptions and numerical techniques employed in modeling stellar adiabatic mass loss in section~\\ref{sec2}, with initial results and their possible application to binary population synthesis in section~\\ref{sec3}. In section~\\ref{sec4}, we offer a short discussion of some  remaining problems. ",
        "conclusions": "\\label{sec4} RLOF process and CE evolution are very important in the formation of hot subdwarf stars and other binary systems. Earlier published studies of \\citet[and references therein]{hjel87} provided useful qualitative insights in RLOF and CE evolution, but left considerable room for improvement, especially with regard to the range of evolved phases that need to be addressed. These deficiencies were in part redressed in later work \\citep{hjel89b}, but only a fragment of that work was ever published \\citep{hjel89a}.  The stellar adiabatic mass loss models described in this paper considerably extend the scope of Hjellming's work.  They allow us not only to study the interior structure of donor stars undergoing dynamical timescale mass transfer, but also to evaluate the stability criteria for dynamical mass transfer in binary population synthesis.  An initial application of our models to $1~M_\\odot$ and $10~M_\\odot$ stars has been described in this paper.  The reader should beware that the these models have limitations to their usefulness, as they depend on being able to separate the process of dynamical relaxation to hydrostatic equilibrium from that of thermal relaxation to thermal equilibrium.  On the main sequence, for example, the dynamical time scale of the Sun is roughly $10^{-10}$ that of its thermal time scale, making the adiabatic mass loss model an excellent description of the asymptotic behavior of a solar-type star to very rapid mass loss.  However, at the opposite extreme in the HR diagram, stars near the tip of the RGB or AGB have thermal time scales approaching their dynamical time scales, and the entire mass transfer process may require full-blown time-dependent modeling. This convergence of thermal and dynamical time scales is broadly related to the abrupt expansion seen in our models of $1~M_\\odot$ and $10~M_\\odot$ stars at the tip of the giant branch (TGB), when their outermost superadiabatic layers are stripped away.  In reality, Roche lobe overflow is far from spherically-symmetric, as treated here (by necessity), but that superadiabatic expansion may nevertheless reflect a real physical phenomenon.  Convection itself is not spherically symmetric, and with convective velocities in the superadiabatic zone approaching sound speed, it may be possible for rising flows near the inner Lagrangian point to bridge the potential to the companion's Roche lobe while the donor still lies well within its own Roche limit.  We are not now able to pass judgement on that possibility.  When dynamical instability occurs, common envelope evolution almost certainly follows, as the thermal time scale of the accreting star is invariably longer than that of the donor, which itself is generally much longer than its dynamical time scale.  This ordering of time scales ensures that the envelope cannot cool efficiently on the transfer time scale, but remains extended and engulfs both stellar cores.  It may happen that thermal time scale mass loss from the donor is still rapid enough to form a common envelope, but this case can only arise if the system is stable against dynamical mass transfer, while unstable to thermal time scale mass transfer. Since the stellar dynamical time scales of both donor and accretor are shorter than their thermal time scales, the prospect arises of a quasistatic common envelope, that is, of formation of a contact binary.  Such large numbers of such objects are known (the W UMa systems, as examples) that they must be very long-lived, evolving in a very different fashion from CE evolution as we have used that term above.  We believe that stellar adiabatic mass loss models provide the most useful approach to date toward defining the limits of dynamical stability in interacting binaries, essential input to the construction of binary population synthesis models."
    },
    "1005/1005.3619_arXiv.txt": {
        "abstract": "We have derived Fe abundances of 16 solar-type Pleiades dwarfs by means of an equivalent width analysis of \\ion{Fe}{1} and \\ion{Fe}{2} lines in high-resolution spectra obtained with the Hobby - Eberly Telescope and High Resolution Spectrograph.  Abundances derived from \\ion{Fe}{2} lines are larger than those derived from \\ion{Fe}{1} lines (herein referred to as over-ionization) for stars with \\teff\\ $< 5400$ K, and the discrepancy (\\dfe\\ = [\\ion{Fe}{2}/H] - [\\ion{Fe}{1}/H]) increases dramatically with decreasing \\teff, reaching over 0.8 dex for the coolest stars of our sample.  The Pleiades joins the open clusters M\\,34, the Hyades, IC\\,2602, and IC\\,2391, and the Ursa Major moving group, demonstrating ostensible over-ionization trends.  The Pleiades \\dfe\\ abundances are correlated with \\ion{Ca}{2} infrared triplet and H$\\alpha$ chromospheric emission indicators and relative differences therein.  Oxygen abundances of our Pleiades sample derived from the high-excitation \\ion{O}{1} triplet have been previously shown to increase with decreasing \\teff, and a comparison with the \\dfe\\ abundances suggests that the over-excitation (larger abundances derived from high excitation lines relative to low excitation lines) and over-ionization effects that have been observed in cool open cluster and disk field main sequence (MS) dwarfs share a common origin.  Curiously, a correlation between the Pleiades \\ion{O}{1} abundances and chromospheric emission indicators does not exist.  Star-to-star \\ion{Fe}{1} abundances have low internal scatter ($< 0.11$ dex), but the abundances of stars with \\teff $< 5400$ K are systematically higher compared to the warmer stars.  The cool star [\\ion{Fe}{1}/H] abundances cannot be connected directly to over-excitation effects, but similarities with the \\dfe\\ and \\ion{O}{1} triplet trends suggest the abundances are dubious.  Using the [\\ion{Fe}{1}/H] abundances of five stars with \\teff\\ $> 5400$ K, we derive a mean Pleiades cluster metallicity of [Fe/H] $= +0.01 \\pm 0.02$. ",
        "introduction": "Studies of Galactic chemical evolution are dependent on accurately derived abundances of stars spanning all ages, populations, kinematics, masses, and metallicities.  Stars with masses $M \\leq 1 \\mathrm{M}_{\\odot}$ are especially important given their dominance of the initial mass function \\citep[IMF; e.g.,][]{2002Sci...295...82K}.  Abundance studies utilizing high-resolution spectroscopy and local thermodynamic equilibrium (LTE) analyses of near-solar metallicity G and K dwarfs in open clusters and in the disk field, however, have revealed that the observed abundances of at least some elements derived for these cool main sequence (MS) dwarfs may be spurious.  In particular, studies have found evidence of over-ionization and over-excitation\\footnotemark[6], i.e., larger abundances are derived from lines of singly ionized species compared to neutral species and from high-excitation lines of neutral species compared to low-excitation lines, respectively. The first indication that some abundances derived for cool MS stars are problematic may have come from \\citet{1974ApJS...27..405O}, who found the over-ionization of Sc, Ti, Cr, and Fe for a sample of 10 K dwarfs ($4800 \\leq$ \\teff $\\leq 5600$ K) in the solar neighborhood.  After a careful analysis of the procedures and stellar parameters used in the abundance derivations, the author was unable to account for the overabundances of the ionized species.  \\citet{1998A&AS..129..237F} found similar over-ionization results for Sc, V, Cr, Fe, and Y in five field K dwarfs ($4510 \\leq$ \\teff\\ $\\leq 4833$ K).  The authors could not exclude an inaccurate temperature scale that is several hundred K too low as a possible source of the anomalous abundances, but in the end, they suggest that non-LTE (NLTE) effects are the more likely cause.  \\footnotetext[6]{In this paper, we use over-ionization and over-excitation to refer to the observed enhanced abundances derived from spectral lines of singly ionized species or from high excitation lines, as opposed to other common usages referring specifically to the non-LTE (NLTE) effects of over-ionization (the mean intensity, $J_{\\nu}$, is larger than the Planck function, $B_{\\nu}$, in lower atomic energy levels), resonance scattering, and photon pumping \\citep[e.g.,][]{1998A&AS..129..237F,2005ARA&A..43..481A}.}   Open clusters have been important to the identification and continued study of the over-excitation/ionization effects, because the presumed internal chemical homogeneity of the clusters provides a baseline with which anomalous abundances can be compared. \\citet{2000ApJ...533..944K} derived O abundances from the high-excitation ($\\chi = 9.15$ eV) near-IR \\ion{O}{1} triplet of a K dwarf in each of the the Pleiades and NGC 2264 open clusters, and in both cases, the abundances were highly enhanced: [O/H] $= +0.85$ and $+0.43$, respectively.  Such high O abundances are not expected for clusters with nominal metallicities of [Fe/H] $= 0.00$ \\citep{1988ApJ...327..389B} and -0.15 \\citep{2000ApJ...533..944K}, respectively.  Following \\citet{2000ApJ...533..944K}, \\citet{2003AJ....125.2085S} derived the abundances of MS dwarfs in the M\\,34 cluster; over-ionization of Fe and over-excitation of Si (the abundances of which were derived from lines with excitation potentials in the range $5.61 \\leq \\chi \\leq 6.19$ eV) are seen in the coolest stars of the sample.  Oxygen abundances of the cool M\\,34 dwarfs, as well as cool Pleiades dwarfs, derived from the high-excitation \\ion{O}{1} triplet are also highly enhanced \\citep{2004ApJ...602L.117S}, confirming the earlier results of \\citet{2000ApJ...533..944K}.  Subsequent to these early open cluster studies, the over-ionization of Fe has been confirmed in the Hyades \\citep{2004ApJ...603..697Y,2006AJ....131.1057S} and Ursa Major (UMa) moving group \\citep{2005PASP..117..911K}, and of Ti in the young pre-MS clusters IC\\,2602 and IC\\,2391 \\citep{2009A&A...501..553D}.  Overabundances of O derived from the \\ion{O}{1} triplet have been reported for the UMa moving group \\citep{2005PASP..117..911K}, the Hyades \\citep{2006ApJ...636..432S}, and IC\\,4665 \\citep{2007ApJ...660..712S}.  Over-excitation effects have been reported for other elements, as well, including S in the Pleiades \\citep {2004ApJ...602L.117S}, Si, Ti, Ni, and Cr in IC\\,4665 \\citep{2005ApJ...635..608S}, Ni in the Hyades \\citep{2006AJ....131.1057S}, and Ca, Ti, and Na in IC\\,2602 and IC\\,2391 \\citep{2009A&A...501..553D}.  Recent abundance analyses of cool field stars have also identified over-excitation/ionization effects \\citep{2004A&A...420..183A,2007A&A...465..271R,2008AJ....135..618C}, confirming the findings of earlier work.  The over-excitation/ionization abundance anomalies are not thought to represent real photospheric overabundances; rather, we believe that they are a signal that our knowledge of cool dwarf atmospheres and/or spectral line formation therein is incomplete.  As of yet, the source or cause of the effects has not been identified.  Systematically erroneous stellar parameters, e.g., an inaccurate \\teff\\ scale, could lead to the observed abundance trends, but in general, unrealistically large parameter errors would have to be present \\citep[e.g.,][]{2005PASP..117..911K,2006ApJ...636..432S}. Furthermore, parameter changes made in response to the overabundances of one element often increase those of another \\citep[e.g.,][]{2003AJ....125.2085S}.  NLTE effects have been suggested as the cause \\citep[e.g.,][]{1998A&AS..129..237F}, but in general, the over-excitation/ionization effects seen in cool dwarfs are in stark contrast to extant NLTE calculations.  For instance, LTE analyses of the high-excitation \\ion{O}{1} triplet in the spectra of MS dwarfs are predicted to result in increasingly discrepant abundances with increasing \\teff\\ for stars with \\teff\\ $> 6000$ K, requiring negative NLTE corrections up to 0.4 -- 0.5 dex at 6500 K for solar metallicity dwarfs \\citep{2003A&A...402..343T, 2009A&A...500.1221F}.  Below 6000 K, the NLTE corrections are predicted to be $< 0.1$ dex and essentially zero below 5500 K \\citep{2003A&A...402..343T}.  Chromospheric emission and photospheric activity (spots, plages, and faculae) have also been suggested sources for the abundance anomalies \\citep[e.g.,][]{2006ApJ...636..432S}.  These inhomogeneities could produce apparent over-excitation/ionization effects within a strict LTE framework.  Continuing our efforts to delineate and understand the observed over-excitation/ionization effects in cool MS dwarfs, we have derived \\ion{Fe}{1} and \\ion{Fe}{2} abundances of 16 Pleiades dwarfs, 15 of which have had O abundances derived from the high-excitation \\ion{O}{1} triplet \\citep{2004ApJ...602L.117S}.  The \\ion{O}{1} abundances evince a steep increase, reaching [O/H] $\\approx 1.0$ dex near 5000 K, and star-to-star dispersion below 5500 K.  We use the newly derived Pleiades Fe abundances to investigate if the over-excitation and over-ionization effects observed in cool MS dwarfs are related and indeed manifestations of the same phenomenon.  Future observational studies that could place stringent constraints on these effects and bring us closer to discovering the source of the anomalous abundances are also discussed. ",
        "conclusions": "We have derived Fe abundances via an EW analysis of \\ion{Fe}{1} and \\ion{Fe}{2} lines in high-resolution and moderate-S/N spectra of 16 MS dwarfs in the Pleiades open cluster. The \\feii\\ abundances increase dramatically relative to \\fei\\ at \\teff\\ below 5400 K, with the difference reaching over 0.8 dex in the coolest stars.  This behavior is akin to what is seen in M\\,34, the Hyades, and the UMa moving group.  Comparison of the \\dfe\\ abundance patterns in the Pleiades and Hyades, as well as the \\otrip\\ abundances in the Pleiades, Hyades, and UMa moving group, suggests that the trends may relax with age, though metallicity may yet prove to be a factor.  Abundances of cool dwarfs in additional open clusters or other stellar associations, especially those older than the Hyades, are needed to determine if either age or metallicity are related to these anomalous abundances.  The \\fei\\ abundances are also higher in Pleiads below 5400 K, but they show no evidence of an increase with decreasing \\teff.  The inability to attribute the high \\fei\\ abundances of the cool stars to the over-excitation effects illustrates the difficulty of quantifying this phenomenon.  With lines of exceptionally high excitation potential such as the \\ion{O}{1} triplet, the over-excitation effect is clearly seen \\citep[e.g.,][]{2006ApJ...636..432S,2007ApJ...660..712S}, but for lines with excitation potentials $\\lesssim 5$ eV, the effect is more difficult to pinpoint.  Our \\ion{Fe}{1} linelist includes transitions ranging in excitation potential from 2.18 to 5.10 eV, but for the individual Pleiades stars, no increase in the line-by-line abundances as a function of excitation potential, like that seen in \\ion{Ni}{1} abundances of cool Hyades dwarfs \\citep[1.83 eV $\\leq \\chi \\leq$ 4.42 eV;][]{2006AJ....131.1057S}, is evident (Figure 4).  Line-to-line sensitivities to the over-excitation/ionization effects have yet to be clearly delineated, and it needs to be determined if there is an excitation potential threshold above which the abundances derived from these lines become enhanced by these effects.  Similarly, it needs to be determined if there is an excitation potential threshold {\\it below} which the opposite occurs, the abundances derived from the low-excitation lines are lower due to these effects.  Such behavior would be expected if the overabundances of high-excitation (singly ionized) lines are due to the overpopulation of high-excitation (singly ionized) electronic states at the expense of depopulating low-excitation states.  Whether or not the over-excitation effects impact the spectroscopic derivation of stellar parameters (\\teff, $\\log g$, and $\\xi$), an approach not adopted here, also needs to be determined.  Future investigations of these effects will require high-quality high-resolution spectroscopy so that accurate line-by-line abundances can be derived, even from features of just a few m{\\AA} in strength.  A strong correlation between the \\dfe\\ and \\otrip\\ abundances of the Pleiades dwarfs is evident in Figure 2, suggesting that the over-excitation/ionization effects share a common cause or origin.  Chromospheric emission and photospheric spots have been shown to be promising culprits, but to this point, the data are inconclusive.  Whereas strong correlations between the Pleiades \\dfe\\ and chromospheric emission indicators and their residuals exist, they do not exist between \\otrip\\ and chromospheric emission.  These contradictory results complicate the interpretation of the observed over-excitation/ionization effects and will have to be addressed by future studies.  Also, comparing abundances and chromospheric emission indicators measured using different spectra may not provide an accurate test of a true correlation because of potential temporal changes in chromospheric activity levels.  Future investigations into the over-excitation/ionization effects in cool open cluster dwarfs should make every effort to derive chromospheric activity levels from the same spectra so that any possible relation between the two can be more definitively delineated.  Determining the influence of photospheric spots on abundance derivations is more arduous. Multicomponent model atmospheres- simulating photospheres with different areal coverages of hot, cool, and quiescent spots- have been shown to able to reproduce the measured EWs of the \\ion{O}{1} triplet in a sample of Hyades stars, but, while such exercises are useful and demonstrate the plausibility of the photospheric spot hypothesis, the results are only suggestive.  Observationally, a simultaneous photometric and spectroscopic monitoring program could be used to identify any correlated changes in spot coverage and spectral line strengths.  Such observational constraints would be helpful to determine if spotted photospheres affect high-resolution abundance derivations.  Despite the challenges, the possible connection between spots, over-excitation/ionization effects, and pre-MS Li depletion should provide sufficient motivation for future efforts.  A final conclusion that can be drawn from this study is that those carrying out spectroscopic abundance analyses of open clusters should heed caution when their samples include cool dwarfs, particularly those with \\teff\\ $< 5400$ K.  Including the abundances of these stars may skew cluster mean abundances.  A similar caution may be needed for those studying cool MS dwarfs in the disk field, as well.  Further investigations into the sensitivity of the over-excitation/ionization effects to excitation potential, first ionization potential, stellar age, and stellar metallicity are needed in order to identify the extent and ubiquity of these effects."
    },
    "1005/1005.0040_arXiv.txt": {
        "abstract": "We discuss the properties of orbits within the influence sphere of a supermassive black hole (BH), in the case that the surrounding star cluster is nonaxisymmetric. There are four major orbit families; one of these, the pyramid orbits, have the interesting property that they can approach arbitrarily closely to the BH. We derive the orbit-averaged equations of motion and show that in the limit of weak triaxiality, the pyramid orbits are integrable: the motion consists of a two-dimensional libration of the major axis of the orbit about the short axis of the triaxial figure, with eccentricity varying as a function of the two orientation angles, and reaching unity at the corners. Because pyramid orbits occupy the lowest angular momentum regions of phase space, they compete with collisional loss cone repopulation and with resonant relaxation in supplying matter to BHs. General relativistic advance of the periapse dominates the precession for sufficiently eccentric orbits, and we show that relativity imposes an upper limit to the eccentricity: roughly the value at which the relativistic precession time is equal to the time for torques to change the angular momentum. We argue that this upper limit to the eccentricity should apply also to evolution driven by resonant relaxation, with potentially important consequences for the rate of extreme-mass-ratio inspirals in low-luminosity galaxies. In giant galaxies, we show that capture of stars on pyramid orbits can dominate the feeding of BHs, at least until such a time as the pyramid orbits are depleted; however this time can be of order a Hubble time. ",
        "introduction": "Following the demonstration that self-consistent equilibria could be constructed for triaxial galaxy models \\citep{Schwarzschild1979,Schwarzschild1982}, observational evidence gradually accumulated for non-axisymmetry on large (kiloparsec) scales in early-type galaxies \\citep{Franx1991,Statler2004,Cappellari2007}. On smaller scales, imaging of the centers of galaxies also revealed a wealth of features in the stellar distribution that are not consistent with axisymmetry, including bars, bars-within-bars, and nuclear spirals \\citep{Shaw1993,ErwinSparke2002,Seth2008}. In the nuclei of low-luminosity galaxies, the non-axisymmetric features may be recent or recurring, associated with ongoing star formation; in luminous elliptical galaxies, central relaxation times are so long that triaxiality, once present, could persist for the age of the universe.  In a triaxial nucleus, torques from the stellar potential can induce gradual changes in the eccentricities of stellar orbits, allowing stars to find their way into the central BH. Gravitational two-body scattering also drives stars into the central BH, but only on a time scale of order the central relaxation time, which can be very long, particularly in the most luminous galaxies. Simple arguments suggest that the feeding of stars to the central BHs in many galaxies is likely to be dominated by large-scale torques rather than by two-body relaxation \\citep[e.g.][]{MerrittPoon2004}.  This paper discusses the character of orbits near a supermassive BH in a triaxial nucleus. The emphasis is on low-angular-momentum, or ``centrophilic,'' orbits, the orbits that come closest to the BH. Self-consistent modelling \\citep{PoonMerritt2004} reveals that a large fraction of the orbits in triaxial BH nuclei can be centrophilic.  Within the BH influence sphere, orbits are nearly Keplerian, and the force from the distributed mass can be treated as a small perturbation which causes the orbital elements (inclination, eccentricity) to change gradually with time. A standard way to deal with such motion is via orbit averaging \\cite[e.g.][]{SandersVerhulst1985}, i.e., averaging the equations of motion over the short time scale associated with the unperturbed Keplerian motion. The result is a set of equations describing the slow evolution of the remaining orbital elements due to the perturbing forces. This approach was followed by \\cite{SridharTouma1999} for motion in an axially-symmetric nucleus containing a massive BH, and by \\cite{SambhusSridhar2000} for motion in a constant-density triaxial nucleus.  In their discussion of motion in triaxial nuclei, \\citet{SambhusSridhar2000} passed over one important class of orbit: the centrophilic orbits, i.e., orbits that pass arbitrarily close to the BH. Examples of centrophiliic orbits include the two-dimensional ``lens'' orbits \\citep{SridharTouma1997,SridharTouma1999} and the three-dimensional ``pyramids'' \\citep{MerrittValluri1999, PoonMerritt2001}. Centrophilic orbits are expected to dominate the supply of stars and stellar remnants to a supermassive BH \\citep[e.g.][]{MerrittPoon2004} and are the focus of the current paper.  The paper is organized as follows. In \\S\\ref{sec_potential} we present a model for the gravitational potential of a triaxial nuclear star cluster, which is more general than that studied in \\cite{SambhusSridhar2000}, but which has many of the same dynamical features. Then in \\S\\ref{sec_orbitavg} we write down the orbit-averaged equations of motion, and in \\S\\ref{sec_analysis} present a detailed analytical study of their solutions, with emphasis on the case where the triaxiality is weak and the eccentricity large. In this limiting case, the averaged equations of motion turn out to be fully integrable. In \\S\\ref{sec_capture} we derive the equations that describe the rate of capture of stars on pyramid orbits by the BH. Comparison of orbit-averaged treatment with real-space motion is made in \\S\\ref{sec_realspace}, to test the applicability of the former. In \\S\\ref{sec_GR} we consider the effect of general relativity on the motion, which imposes an effective upper limit on the eccentricty. \\S\\ref{sec_RR} discusses the connection with resonant relaxation: we argue that a similar upper limit to the eccentricity should characterize orbital evolution in the case of resonant relaxation. Finally, in \\S\\ref{sec_estimates} we make some quantitative estimates of the importance of pyramid orbits for capture of stars in galactic nuclei. \\S\\ref{sec_summary} sums up. ",
        "conclusions": "\\label{sec_summary}  We discussed the character of orbits within the radius of influence $\\rh$ of a supermassive BH at the center of a triaxial star cluster. The motion can be described as a perturbation of Keplerian motion; we derive the orbit-averaged equations and explore their solutions both analytically (when the triaxiality is small) and numerically. Orbits are found to be mainly regular in this region. There exist three families of tube orbits; a fourth orbital family, the pyramids, can be described as eccentric Keplerian ellipses that librate in two directions about  the short axis of the triaxial figure. At the ``corners'' of the pyramid, the angular momentum reaches zero, which means that stars on these orbits can be captured by the BH. We derive expressions for the rate at which stars on pyramid orbits would be lost to the BH; there are many similarities with the more standard case of diffusional loss cone refilling, but also some important differences, due  to the fact that the approach to the loss cone is deterministic for the pyramids, rather than statistical. The inclusion of general relativistic precession is shown to impose a lower bound on the angular momentum. We argue that a similar lower bound should  apply to orbital evolution in the case that the torques are due to resonant relaxation. The rate of consumption of stars from pyramid orbits is likely to be substantially greater than the rate due to two-body relaxation in the most luminous galaxies, although in the absence of mechanisms for orbital repopulation, these high consumption rates would  only be maintained until such a time as the pyramid orbits have been drained; however the latter time can be measured in billions of years."
    },
    "1005/1005.2860_arXiv.txt": {
        "abstract": "We report the detections of two substellar companions orbiting around evolved intermediate-mass stars from precise Doppler measurements at Subaru Telescope and Okayama Astrophysical Observatory. HD 145457 is a K0 giant with a mass of 1.9 $M_{\\odot}$ and has a planet of minimum mass $m_2\\sin i=2.9 M_{\\rm J}$ orbiting with period of $P=176$ d and eccentricity of $e=0.11$. HD 180314 is also a K0 giant with 2.6 $M_{\\odot}$ and hosts a substellar companion of $m_2\\sin i=22 M_{\\rm J}$, which falls in brown-dwarf mass regime, in an orbit with $P=396$ d and $e=0.26$. HD 145457 b is one of the innermost planets and HD 180314 b is the seventh candidate of brown-dwarf-mass companion found around intermediate-mass evolved stars. ",
        "introduction": "Over 400 exoplanets have been discovered by various techniques during the past 15 years\\footnote {See, e.g., table at http://exoplanet.eu/}. Among the techniques, precise Doppler technique has been the most powerful method for planet detection around various types of stars including solar-type stars, evolved giants and subgiants, early-type stars, and so on (e.g. Udry \\& Santos 2007 and references therein). Recently, precise Doppler measurements in infrared wavelength region started to explore planets around very low-mass stars down to 0.1 $M_{\\odot}$ (Bean et al. 2010). These surveys will help us to understand properties of planets as a function of stellar mass, age, evolutionary stage, and so on, and thus provide us more general picture of planet formation and evolution.  In particular, planets around giants and subgiants have been extensively surveyed over the past several years mainly from the viewpoint of planet searches around intermediate-mass (1.5--5$M_{\\odot}$) stars (e.g. Frink et al. 2002; Setiawan et al. 2005; Sato et al. 2008b; Hatzes et al. 2005, 2006; Johnson et al. 2010; Lovis \\& Mayor 2007; Niedzielski et al. 2009b; D$\\ddot{\\rm{o}}$llinger et al. 2009; de Medeiros et al. 2009; Liu et al. 2009; Omiya et al. 2009). Intermediate-mass dwarfs, namely BA-type dwarfs, are more difficult for Doppler planet searches because of paucity of spectral features and their rotational broadening. It is thus difficult to achieve high measurement precision in radial velocity (but see eg. Galland et al. 2005, 2006). On the other hand, those in evolved stages, namely GK-type giants and subgiants, have many sharp absorption lines in their spectra suitable for precise radial velocity measurements, which makes them promising targets for Doppler planet searches. Actually, more than 30 substellar companions to such evolved stars have been already found and they have shown remarkable properties: more than twice higher occurrence rate of giant planets for intermediate-mass subgiants than that for lower-mass stars (Johnson et al. 2007b; Bowler et al. 2010), larger typical mass of giant planets (Lovis et al. 2007; Omiya et al. 2009), lack of inner planets with semimajor axes $<$0.6 AU (Johnson et al. 2007a; Sato et al. 2008a; Niedzielski  et al. 2009a), and lack of metal-rich tendency in host stars (Pasquini et al. 2007; Takeda et al. 2008). All these statistical properties should be confirmed by collecting a larger number of samples.  Since 2001, we have been carrying out a Doppler planet search program targeting 300 GK giants at Okayama Astrophysical Observatory (OAO) and have discovered 9 planets and 1 brown dwarf so far from the program (Sato et al. 2003, 2007, 2008ab; Liu et al. 2008). In order to further extend the survey, we have established an international consortium between Chinese, Korean, and Japanese researchers using 2m class telescopes in the three countries (East-Asian Planet Search Network; Izumiura 2005), which recently announced discoveries of a planet (Liu et al. 2009) and a brown dwarf (Omiya et al. 2009) around GK giants. A total of about 600 GK giants are now under survey by the consortium.  In this paper, we report the detections of two new substellar companions around intermediate-mass giants (HD 145457 and HD 180314) from our newly started planet search program using Subaru 8.2m telescope and the above 2m class telescopes. The substellar companions presented here were uncovered by the initial screening with Subaru and followed up with OAO 1.88m telescope. We describe the outline of the Subaru survey in Section 2 and the observations in Section 3. The stellar properties are presented in Section 4, and radial velocities, orbital solutions, and results of line shape analysis are provided in Section 5. Section 6 is devoted to summary and discussion. ",
        "conclusions": "We here reported the two new substellar companions to K0 III giants from Subaru and OAO planet search programs. The discoveries add to the recent growing population of substellar companions around evolved intermediate-mass stars.  HD 145457 b ($m_2\\sin i=2.9M_{\\rm J}$, $a=0.76$ AU) is one of the innermost planets found around giants. All the planets currently known around evolved intermediate-mass stars orbit at a$\\ge$0.6 AU (Johnson et al. 2007a; Sato et al. 2008a; Niedzielski  et al. 2009a). Two possible causes of the lack of inner planets has been proposed: they are originally deficient or engulfed by the central stars. In the case of intermediate-mass subgiants  ($<2M_{\\odot}$), the former scenario is considered to be appropriate because they are obviously less evolved and have stellar radii of $\\lesssim 6R_{\\odot}$, which means that they could only have engulfed very short-period planets like hot-Jupiters (e.g. Johnson et al. 2007a). It is not easy, however, to discriminate between the two scenarios in the case of intermediate-mass giants (typically $\\ge2M_{\\odot}$) because it is difficult to know accurate evolutionary status of giants. The inner planets could be tidally engulfed by the central stars at the phase of the tip of red giant branch (RGB). Thus if we observe giants that have already passed through the tip of RGB, namely core-helium burning stars, we could not find any inner planets around the stars even if the planets originally existed (Sato et al. 2008a; Villaver \\& Livio 2009). However, it is apparently difficult to distinguish such giants from those ascending RGB for the first time because they locate at nearly the same region on the HR diagram (red clump region). Based on stellar evolutionary models (cf.evolutionary tracks by Girardi et al. 2000), core-helium burning giants stay at the clump region for $\\sim$100 times longer than first ascending RGB stars do. Therefore, if we can collect hundreds of planets around such giants, it will become possible to verify the causes of the lack of inner planets statistically.  HD 180314 b has a minimum mass of 22 $M_{\\rm J}$ and orbits around the central star with 2.6$M_{\\odot}$. This is the 7 th brown-dwarf-mass (13--80$M_{\\rm J}$) companion to evolved intermediate-mass stars (Hatzes et al. 2005; Lovis \\& Mayor 2007; Liu et al. 2008; Omiya et al. 2009; Niedzielski et al. 2009b\\footnote{BD+20 2457 has two brown-dwarf-mass companions in the system.}). Actually all of their host stars are estimated to be more massive than 2.5 $M_{\\odot}$, suggesting that more massive stars tend to have more massive companions (Lovis \\& Mayor 2007; Omiya et al. 2009). Several scenarios have been proposed for the formation of brown-dwarf-mass companions including gravitational collapse in protostellar clouds like stellar binary systems (Bonnell \\& Bastien 1992; Bate 2000) and gravitational instability in protostellar disks (Boss 2000; Rice et al. 2003). Even by core-accretion scenario in protoplanetary disks, super-massive companions with $\\gtsim10M_{\\rm J}$ could be formed on a certain truncation condition for gas accretion (Ida \\& Lin 2004; Alibert et al. 2004; Mordasini et al. 2007). If the companions form like stellar binary systems, they are expected to have a wide variety of orbital eccentricity. The above 7 brown dwarf candidates, however, have relatively low eccentricies of 0--0.3, which may be favored by the scenarios that they formed in circumstellar disks and have not experienced significant gravitational interaction with other companions. One more thing to note here is that metallicity of the host stars of the above brown dwarf candidates are ranging from $\\rm {[Fe/H]}=-1$ to 0.2. Although the number of samples is still small, this suggests that the formation mechanism is independent of or less sensitive to metallicity in this range. \\\\  This research is based on data collected at Subaru Telescope and Okayama Astrophysical Observatory (OAO), which are operated by National Astronomical Observatory of Japan (NAOJ). We are grateful to all the staff members of Subaru and OAO for their support during the observations. We thank the National Institute of Information and Communications Technology for their support on high-speed network connection for data transfer and analysis. BS is supported by MEXT's program \"Promotion of Environmental Improvement for Independence of Young Researchers\" under the Special Coordination Funds for Promoting Science and Technology, and by Grant-in-Aid for Young Scientists (B) No.20740101 from the Japan Society for the Promotion of Science (JSPS). YJL is supported by the National Natural Science Foundation of China under grant number 10803010 and 10821061. BCL acknowledge the Astrophysical Research Center for the Structure and Evolution of the Cosmos (ARCSEC, Sejong University) of the Korea Science and Engineering Foundation (KOSEF) through the Science Research Center (SRC) program. This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France."
    },
    "1005/1005.0923_arXiv.txt": {
        "abstract": "The circumgalactic medium (CGM) around galaxies is believed to record various forms of galaxy feedback and contain a significant portion of the ``missing baryons'' of individual dark matter halos. However, clear observational evidence for the existence of the hot CGM is still absent. We use intervening galaxies along 12 background AGNs as tracers to search for X-ray absorption lines produced in the corresponding CGM. Stacking \\chandra grating observations with respect to galaxy groups and different luminosities of these intervening galaxies, we obtain spectra with signal-to-noise ratios of 46-72 per 20-m\\AA\\ spectral bin at the expected \\ovii\\ K$\\alpha$ line. We find no detectable absorption lines of \\cvi, \\nvii, \\ovii, \\oviii, or \\neix. The high spectral quality allows us to tightly constrain upper limits to the corresponding ionic column densities (in particular $\\log[N{\\rm_{OVII}(cm^{-2})}]\\le$14.2--14.8). These nondetections are inconsistent with the Local Group hypothesis of the X-ray absorption lines at $z\\simeq0$ commonly observed in the spectra of AGNs. These results indicate that the putative CGM in the temperature range of $10^{5.5}-10^{6.3}$ K may not be able to account for the missing baryons unless the metallicity is less than 10\\% solar. ",
        "introduction": "\\label{sec:intro}  Modern simulations are converging on the formation and evolution of the dark galactic halos as well as the large-scale intergalactic structure of the Universe (e.g., \\citealt{nav96, she01, cen99, dav99, spr05}). However, serious difficulties are present in reproducing the visible parts of galaxies, including the mass/luminosity functions at low and high masses, the angular momentum distribution of galactic disks, and the supermassive black hole (SMBH) and bulge-mass relationship (e.g., \\citealt{gil08, pri09, cat09}). For instance, the ratio of the visible baryonic mass to the gravitational mass of all galaxies is 3-10 times less than the cosmic value ($0.167\\pm0.006$) inferred from the {\\sl Wilkinson Microwave Anisotropy Probe} (\\citealt{kom09}; also see \\citealt{bre09} and references therein); the discrepancy seems to be more severe in less massive galaxies than in more massive ones (e.g., \\citealt{hoe05, mcg08}). This missing link between the intergalactic medium (IGM), dark halos, and visible galaxies is attributed largely to our poor understanding of the complex ``gastrophysics'', in particular the coupling between gas and galaxy feedback (e.g., \\citealt{dav07}).  Galaxy feedback comes in many forms. Starbursts are known to generate galactic winds, which must be chiefly responsible for the chemical enrichment and non-gravitational heating of the IGM at high redshifts (e.g., \\citealt{mac99}). Nuclear starbursts are also believed to be intimately related to the formation and growth of the SMBHs, which are an additional source of feedback, especially in massive galaxies (e.g., \\citealt{str02, swa06}). While the radiation pressure in a high-accretion phase can be a key driving force of the winds, the mechanical energy injection in a radiatively inefficient accretion phase may balance the cooling of the surrounding gas. Feedback from stars, even from evolved ones alone (e.g., Type Ia supernovae), can play an essential role in maintaining hot gaseous halos around galaxies \\citep{tang09a, tang09b}. Cosmic-ray pressure can also play an important role in driving large-scale galactic outflows \\citep{eve08}. The relative importance of these various forms of the galaxy feedback, however, remains very unclear. Direct observational constraints on the physical process and history of the galaxy feedback are thus badly needed.  A potentially effective approach is to observe the circumgalactic medium (CGM) around nearby galaxies. On scales from a few tens of kpc up to $\\sim 1$ Mpc, the CGM includes the gas that has been significantly affected by the galaxy feedback but is outside of the boundaries of the traditionally-known stellar and multi-phase interstellar components of such galaxies. It is quite possible that the CGM around the galaxies contains their ``missing'' baryons (e.g., \\citealt{deh98,som06}).  Searches for the CGM have been conducted in ultraviolet (UV), but the association between galaxies and IGM absorbers still remains an unsettled issue. Within impact distances of $r_\\rho\\lsim350$ kpc from the sight lines of background QSOs, all galaxies with $L>0.1L^*$ seem to have ``associated'' Ly$\\alpha$ or other low ionized absorbers (e.g., \\citealt{chen08,wak09}). However, not all Ly$\\alpha$ absorbers have an associated galaxy found within 1$h^{-1}$ Mpc (e.g., \\citealt{mor93,sto06}). These absorbers could be either associated with faint galaxies that are beyond the detection limits of current galaxy surveys or not related to any galaxies at all (some absorbers are detected in voids; \\citealt{chen09, sto07}.) Also, it is possible that much of the injection into the IGM occurred at high redshifts, and in the time that has subsequently passed, substantial separations in space and velocity have developed between the ejected matter and the galaxies from which it originated \\citep{tri06}. The highly ionized absorbers, \\ovi\\ absorbers in particular, may trace the CGM at high temperatures. Covering fraction of these absorbers varies with respect to galaxy luminosities and the impact distances of these absorbers to their nearest galaxies \\citep{sto06, pro06, gan08, wak09}.  X-ray observations could provide complementary or even key constraints on the existence and properties of the CGM. The bulk of the CGM, at least for galaxies with masses similar to and higher than our Galaxy, is believed to be in a gaseous phase at temperatures of $T\\gsim 10^6$~K and mainly emits and absorbs X-ray photons (e.g., \\citealt{bir03}). However, because of the low density and potentially low metallicity of the CGM, its X-ray emission, except for that in galaxy clusters where the missing baryon problem becomes less severe, is expected to be hard to detect (e.g., \\citealt{han06}), and the observational evidence is still absent (e.g., \\citealt{ras09, and10}).  The high spectral resolution grating instruments aboard \\chandra\\ and \\xmm\\ now provide us with a new powerful tool --- X-ray absorption line spectroscopy --- to probe the low-surface brightness diffuse hot CGM. Unlike X-ray emission that is sensitive to the emission measure of the hot gas, absorption lines produced by helium- and hydrogen-like ions such as \\ovii, \\oviii, and \\neix\\ directly probe their column densities, which are proportional to the mass of the gas and are sensitive to its thermal, chemical, and kinematic properties. Indeed, both the \\chandra\\ and \\xmm\\ grating observations of bright AGNs and X-ray binaries have been used to firmly detect and characterize the global hot gas in and around the Milky Way (e.g., \\citealt{wang05, yao05, yao06, yao07, fang06, bre07}). However, our location in the midst of the hot gas makes it hard to determine the presence of the CGM around the Galaxy; our previous investigations only yield tentative upper limits to the column densities of the CGM beyond the hot gas-rich Galactic disk (e.g., \\citealt{yao08, yao09a}).  In this work, we search for absorption lines of \\cvi, \\nvii, \\ovii, \\oviii, and \\neix\\ produced in the CGM of intervening galaxies along 12 AGN sight lines by extensively exploring \\chandra\\ grating observations. Because detection sensitivities of current high resolution X-ray instruments are still limited, to enhance the signal-to-noise (S/N) ratio, we stack spectra with respect to group membership and luminosity properties of intervening galaxies. This stacking technique, as implemented in our previous paper \\citep{yao09b}, enables us to probe the absorbing gas to unprecedented low column densities.  The paper is organized as follows. In Section~\\ref{sec:data}, we discuss the background AGNs and survey of intervening galaxies, and we describe the data reduction process. We search for and measure the X-ray absorption lines produced in the CGM in Section~\\ref{sec:results}, and discuss the implications of our results % in Section~\\ref{sec:dis}.  Throughout the paper, we adopt the Schechter luminosity function with characteristic luminosity $L_B^*\\simeq2\\times10^{10}L_\\odot$ (or absolute B magnitude $M_B^*\\simeq-20.5$) and use the cosmology with $H_0=73~{\\rm km~s^{-1}~Mpc^{-1}}$, $\\Omega_{M}=0.3$, and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": ""
    },
    "1005/1005.0112_arXiv.txt": {
        "abstract": "Many astrophysical phenomena are highly subsonic, requiring specialized numerical methods suitable for long-time integration.  In a series of earlier papers we described the development of {\\tt MAESTRO}, a low Mach number stellar hydrodynamics code that can be used to simulate long-time, low-speed flows that would be prohibitively expensive to model using traditional compressible codes.  {\\tt MAESTRO} is based on an equation set derived using low Mach number asymptotics; this equation set does not explicitly track acoustic waves and thus allows a significant increase in the time step. {\\tt MAESTRO} is suitable for two- and three-dimensional local atmospheric flows as well as three-dimensional full-star flows.  Here, we continue the development of {\\tt MAESTRO} by incorporating adaptive mesh refinement (AMR). The primary difference between {\\tt MAESTRO} and other structured grid AMR approaches for incompressible and low Mach number flows is the presence of the time-dependent base state, whose evolution is coupled to the evolution of the full solution.  We also describe how to incorporate the expansion of the base state for full-star flows, which involves a novel mapping technique between the one-dimensional base state and the Cartesian grid, as well as a number of overall improvements to the algorithm.  We examine the efficiency and accuracy of our adaptive code, and demonstrate that it is suitable for further study of our initial scientific application, the convective phase of Type Ia supernovae. ",
        "introduction": "Many astrophysical phenomena of interest occur in the low Mach number regime, where the characteristic fluid velocity is small compared to the speed of sound.  Some well-known examples are the convective phase of Type Ia supernovae (SN~Ia)~\\citep{hoflichstein:2002,kuhlen-ignition:2005,ZABNW:IV}, classical novae~\\citep{glasner:2007}, convection in stars~\\citep{meakin:2007}, and Type I X-ray bursts~\\citep{Lin:2006}. Such problems require a numerical approach capable of resolving phenomena over time scales much longer than the characteristic time required for an acoustic wave to propagate across the computational domain.  In a series of papers (see \\citet{ABRZ:I}---henceforth Paper I, \\citet{ABRZ:II}---henceforth Paper II, \\citet{ABNZ:III}---henceforth Paper III, and \\citet{ZABNW:IV}---henceforth Paper IV), we have described the initial development of {\\tt MAESTRO}, a low Mach number hydrodynamics code for computing stellar flows using a time step constraint based on the fluid velocity rather than the sound speed. {\\tt MAESTRO} is suitable for two- and three-dimensional local atmospheric flows as well as three-dimensional full-star flows.  All simulations are performed in a Cartesian grid framework, but rely on the presence of a one-dimensional radial base state that describes the average state of the star or atmosphere.  Starting with the development of the low Mach number equation set (see Paper I), we demonstrated how to capture the expansion of the base state in a local atmospheric simulation in response to large-scale heating (Paper II) and incorporate reaction networks (Paper III).  In Paper IV, we presented the initial application of {\\tt MAESTRO}, following the last two hours of convection inside a white dwarf leading up to the ignition of a SN Ia using a three-dimensional, full-star simulation with a base state that is constant in time.  In general, astrophysical flows are highly turbulent.  In the case of the convective period preceding a SN Ia explosion, the Reynolds number is $\\mathcal{O}(10^{14})$ \\citep{Woosley:2004}, far larger than can be modeled on today's supercomputers. Nevertheless, to understand the role of turbulence in these events, we must use increasingly more accurate simulations. In this paper we describe how to incorporate adaptive mesh refinement (AMR), in which we locally refine the Cartesian grid in regions of interest, to allow us to efficiently push to higher spatial resolutions and better capture the turbulent flow in critical regions of the simulation.  The primary difference between {\\tt MAESTRO} and other structured grid AMR approaches for incompressible and low Mach number flows is the presence of the time-dependent base state, whose evolution is coupled to the evolution of the full solution. We also describe how to incorporate a time-dependent base state for full-star problems, which involves a novel mapping technique between the one-dimensional base state and the Cartesian grid. This allows us to properly capture the effects of an expanding base state in full-star simulations. We have also made a number of overall improvements to the algorithm, and all together, these enhancements will allow us to compute more efficient and accurate solutions for our target applications, including the convective phase of SNe Ia and Type I X-ray bursts.  This paper is divided into several sections, along with three detailed appendices.  In \\S \\ref{Sec:Governing Equations}, we present the governing equations. In \\S \\ref{Sec:Overview}, we give an overview of the methodology, referring the reader to the appendices for full details.  In \\S \\ref{Sec:Mapping}, we describe the new mapping procedure between one-dimensional and three-dimensional data structures in full-star problems.  In \\S \\ref{Sec:AMR}, we discuss the extension of the algorithm to include AMR.  In \\S \\ref{Sec:Test Problems}, we describe the results of our test problems.  We conclude with \\S \\ref{Sec:Conclusions}, which includes future plans for scientific investigation. ",
        "conclusions": "\\label{Sec:Conclusions} We have developed a low Mach number hydrodynamics algorithm suitable for full star flows and local atmospheric regions with a time-evolving base state within an AMR framework.  In forthcoming papers, we will use {\\tt MAESTRO} to further our scientific investigation of the convective phase of SNe~Ia. Our previous simulations in Paper IV were at modest resolution and assumed a constant base state.  We are now performing simulations at higher effective resolutions with the use of AMR along with a time-varying base state. As part of this study, we are examining the tagging conditions necessary to model a full-star up to the point of ignition. We are also studying Type I X-ray bursts \\citep{STRO_BILD06,Lin:2006}, which are believed to be caused by the thermonuclear explosive burning of hydrogen and/or helium gas accreted into a thin shell on the surface of neutron stars.  We pose this problem in planar geometry, model a small patch of the star, and refine near the base of the accreted layer."
    },
    "1005/1005.4507_arXiv.txt": {
        "abstract": "Swift opened up a new era in the study of gamma-ray burst sources (GRB). Among a variety of discoveries made possible by Swift, here we focus on GRB\\,090423, the event at z=8.2 which currently holds the record of the most distant celestial object ever caught by human instrumentation. This GRB allowed us to have a direct look at the early Universe. The central engine activity giving origin to the GRB emission  is also discussed starting from the observational findings of an updated GRB X-ray flares catalog. ",
        "introduction": "Each time we build and use innovative state of the art instrumentation we expect to detect new phenomena and details that allows a better understanding of the physics of the objects we are observing. In the case of the Swift mission\\cite{Gehrels04} we were able to build on the experience obtained by Beppo-Sax and on the theoretical developments that helped the design allowing targeting critical points of the GRB physics. The result was a multi-wavelength mission with very fast pointing, accurate astrometry of the transients and real time analysis and communication of the detected targets.\\\\  Out of the variegated science and related high lights we have recently published, here we will discuss GRB\\,090423, a burst holding the high z record for any celestial object so far discovered, and   illustrate some of the results we obtained with the new sample of GRB flares. In other words we will look into the early Universe and perhaps into the activity intimately related to the central engine.\\\\  Nowadays we have evidence that the early star formation, pop III, is followed by the re-ionization (of HI) phase that likely terminated at $z \\sim 6$. As it is well known even a small amount of neutral hydrogen would heavily absorbs at wavelengths smaller than Ly$_{\\alpha}$(Gunn - Peterson effect) so that this region in which some neutral hydrogen still exist in the IGM is easily recognizable and indeed it is a powerful tool to track down the various phases of the process as a function of redshift.\\\\ In the X-ray light curve of the afterglow (generally covering the temporal range of few hundreds second since trigger to a million second after that) we define as flare an event that manifests itself as a sudden change in the observed flux and lasting a time largely smaller than the light curve of the GRB itself. The cause of these later injections of energy is not yet known and while it may be caused either by late shocks among shells emitted during the prompt emission or, more likely, by renewed activity of the central engine, they are certainly not due to external shock (it can be shown that the pulse width divided by the time of occurrence of the pulse is too small to be explained as due to external shock). But no matter what, a clear understanding of the observed emission will lead to the understanding of the mechanism at work and eventually lead to the understanding of the central engine. Any correlation among parameters of flares occurring at different time in different bursts must relate to a behavior of the central engine and may lead to its partial understanding. ",
        "conclusions": "Recent observations challenge the fireball internal shock model. Optical observations do not show strong evidence of the reverse shock\\cite{Kumar07}, and in any case even if present does not seem to be one of the main component of the early optical emission\\cite{Oates09}. The reverse shock is weak or suppressed in the magneto hydro dynamical models of GRBs\\cite{Thompson94,Spruit01,Lyutikov03} so that the goal is to find way to estimate whether $\\sigma_{0}=\\frac{F_{p}}{F_{b}}=\\frac{B_{0}^{2}}{(4\\,\\pi\\,\\gamma_{0}\\,\\rho\\,c^{2})}$ is $\\gg 1$ (the jet is dominated by a Pointing flux) or $\\ll 1$ (the kinetic energy of the baryonic jet dominates).\\\\  \\begin{figure} \\begin{center} \\includegraphics[width=\\hsize]{variability.ps} \\caption{The relative variability of the flares as a function of the width to time ratio has been plotted. The limits shown by the various lines have been calculated according to the equations given in Ref.\\,\\protect\\refcite{Ioka05}. The 4 black dots refer to short GRBs.}\\label{fig:variability} \\end{center} \\end{figure}  We should add incidentally, that the mechanism by which strong magnetic fields and the acceleration of particles to very high energy is achieved is not yet fully known. It needs to be demonstrated that the linear Fermi mechanism is capable of doing the job following the back and for transit of the particles in the shock region.\\\\  To summarize part of the results in a diagram we refer to plot of the ratio $\\Delta F/F$ versus $\\Delta t/t$, Figure\\,\\ref{fig:variability}. Here $\\Delta F$ is the variation of the flux above the underlying continuum over the flux observed in the underlying continuum and $\\Delta t/t$ the flare width over the time of the peak of the flare.\\\\  The complete absence of flares with $\\Delta t/t > 1$ confirms that flares cannot be due to patchy shells. Flares furthermore seem to have an high probability, as expected, to be observed off axis and a large number of flares would, according to this plot, agree with the refreshed shock model.\\\\  Flares are not necessarily the result of late central engine activity, but may be produced in the decelerating phase of the flow. High $\\sigma_{0}$ values may lead to MHD instabilities during the interaction with the interstellar medium. In this case the isotropic equivalent energy emitted in a single flare and produced by a single reconnection event, is related to the ratio width/$t_{peak}$ by the relation:  \\begin{equation} E_{flare} \\leq 5\\,\\epsilon \\left(\\frac{width}{t_{peak}}\\right)^{3} \\frac{E_{forward\\,shock}}{\\alpha^{2}} \\end{equation}  where $\\epsilon \\sim 0.1$ and $\\alpha$ in the range 2 - 4 depending on the density of the ISM. The flares we observe are not in contradiction within errors with this relation but, at the same time, they do not prove it.\\\\ In conclusion we find that flares seem to retain memory of the previous event so that, as time progresses, each flare is weaker and softer of the preceding one. And finally while we are making significant progress in characterizing these events observationally, at the moment there is no satisfactory model explaining their origin, evolution and energetics."
    },
    "1005/1005.4677_arXiv.txt": {
        "abstract": "Based on our recent work on tidal tails of star clusters \\citep{Kuepper09} we investigate star clusters of a few $10^4 \\msun$ by means of velocity dispersion profiles and surface density profiles. We use a comprehensive set of $N$-body computations of star clusters on various orbits within a realistic tidal field to study the evolution of these profiles with time, and ongoing cluster dissolution.\\\\ From the velocity dispersion profiles we find that the population of potential escapers, i.e. energetically unbound stars inside the Jacobi radius,  dominates clusters at radii above about 50\\% of the Jacobi radius. Beyond 70\\% of the Jacobi radius nearly all stars are energetically unbound. The velocity dispersion therefore significantly deviates from the predictions of simple equilibrium models in this regime. We furthermore argue that for this reason this part of a cluster cannot be used to detect a dark matter halo or deviations from Newtonian gravity.\\\\ By fitting templates to the about $10^4$ computed surface density profiles we estimate the accuracy which can be achieved in reconstructing the Jacobi radius of a cluster in this way. We find that the template of King (1962) works well for extended clusters on nearly circular orbits, but shows significant flaws in the case of eccentric cluster orbits. This we fix by extending this template with 3 more free parameters. Our template can reconstruct the tidal radius over all fitted ranges with an accuracy of about 10\\%, and is especially useful in the case of  cluster data with a wide radial coverage  and for clusters showing significant extra-tidal stellar populations. No other template that we have tried can yield comparable results over this range of cluster conditions. All templates fail to reconstruct tidal parameters of concentrated clusters, however.\\\\ Moreover, we find that the bulk of a cluster adjusts to the mean tidal field which it experiences and not to the tidal field at perigalacticon as has often been assumed in other investigations, i.e. a fitted tidal radius is a cluster's time average mean tidal radius and not its perigalactic one.\\\\ Furthermore, we study the tidal debris in the vicinity of the clusters and find it to be well represented by a power-law with a slope of -4 to -5. This steep slope we ascribe to the epicyclic motion of escaped stars in the tidal tails. Star clusters close to apogalacticon show a significantly shallower slope of up to -1, however. We suggest that  clusters at apogalacticon can be identified by measuring this slope. ",
        "introduction": "Velocity dispersion profiles and surface density profiles are among the most basic tools for investigating the structure of star clusters. However, such investigations indicate that the region around the tidal radius, at which the internal acceleration of a star cluster is similar to the tidal acceleration due to the galactic tidal field, is particularly poorly understood.\\\\  Velocity dispersion profiles sometimes show peculiarities which have been discussed in the literature.  \\citet{Drukier98}, for example, observed a flattening in the outer parts of the velocity dispersion profile of the Galactic globular cluster M15 which they interpreted as an effect of tidal heating by the general Galactic tide or by tidal shocks. \\citet{Scarpa03} also found a flattening of the velocity dispersion profile for $\\omega$ Cen and more recently for other Galactic globular clusters like NGC6171, NGC7099 and NGC288 \\citep{Scarpa07}. The deviation from an expected Keplerian fall-off in the velocity dispersion profile occurred in all clusters at radii where the internal gravitational acceleration is about $a_0 = 1.2 \\times 10^{-8}\\mbox{cm s}^{-2}$ and was therefore interpreted by Scarpa et al.~as a hint of Modified Newtonian Dynamics \\citep{Milgrom83} in globular clusters. Alternatively, they briefly discussed the possible effect of tidal heating or a dark matter halo on the cluster stars.\\\\\\  On the contrary, \\citet{McLaughlin03b} found that by fitting a \\citet{Wilson75} model, which has a less sharp cut-off at the tidal radius than the commonly used \\citet{King66} model, to the re-analyzed $\\omega$ Cen data, its velocity dispersion profile could be explained without modifying Newtonian gravity and without adding dark matter. Similar investigations of cluster profiles, e.g. \\citet{Lane09, Lane10} and \\citet{Baumgardt10}, were also not in favour of MOND.\\\\  But it is not only the velocity dispersion profiles of star clusters which behave strangely at the tidal boundary;  their surface density profiles also show peculiarities and sometimes controversial behaviour.\\\\  \\citet{King62} showed that the surface density profiles of many globular clusters can be fitted by a simple analytical formula having a sharp cut-off radius, which could be interpreted as the tidal radius of the cluster. Later he derived a set of physically motivated models, with a similar cut-off radius corresponding to an energy cut-off in the energy distribution function of the cluster stars, which provided an even better fit to cluster profiles \\citep{King66}.\\\\  In contrast to that, \\citet{Elson87}, and more recently \\citet{Gouliermis10}, found that young massive clusters in the LMC show an exponential surface density profile without any tidal truncation at the expected tidal radius. They interpreted these findings as being due to tidal debris which was expelled at birth from the clusters and has not had time to disperse yet.\\\\  Furthermore, \\citet{Cote02}, \\citet{Carraro07} and \\citet{Carraro09} find the outer halo Milky-Way globular clusters Palomar 13, Whiting 1 and AM 4, respectively, to have a significant excess of stars at the tidal boundary, which makes any fit to the surface density data very inaccurate. For all clusters they find the radial surface density profile, $\\Sigma(R)$, to be well represented by a power-law $R^\\eta$ with slopes of about $\\eta \\simeq -1.8$. This excess of stars is interpreted by the authors as heavy mass loss in a final stage of dissolution. The same was found for Palomar 5, which is a well studied MW globular cluster close to the apogalacticon of its orbit \\citep{Odenkirchen03}.\\\\  Moreover, \\citet{McLaughlin05} showed that most globular clusters of the Milky Way, the LMC and SMC, as well as of the Fornax dwarf spheroidal, are more extended than can be explained by \\citet{King66} models, and therefore are better represented by \\citet{Wilson75} models. In this context they emphasised the lack of physical understanding of this phenomenon.\\\\  On the contrary, \\citet{Barmby09} found that of 23 young massive clusters in M31 most were better fitted by \\citet{King66} models as these clusters do not show extended haloes.\\\\  As a consequence of this lack of understanding, \\citet{Barmby09} asked in their investigation how robust the physical parameters are which were derived in such analyses. This they tried to estimate by analyzing artificial clusters in the same way as the real observations. A similar analysis has also been performed by \\citet{Bonatto08}, although they tested the analytic formula of \\citet{King62} with artificial observations under various limiting conditions. Both investigations came to the conclusion that physical parameters can in principle be well recovered from such idealized mock observations.\\\\  But probably the idealized nature of these investigations is misleading, since both tests were performed with dynamically unevolved clusters. However, a self-consistent test of deriving physical parameters from a set of numerical computations of star clusters with a range of initial parameters has not been performed yet. Just a few investigations have touched this topic so far by means of numerical computations (e.g. \\citealt{Capuzzo05}, \\citealt{Drukier07}, \\citealt{Trenti10}). This is due to the fact that fast codes for globular cluster integration like Fokker-Planck or Monte-Carlo codes are not able to address this problem properly, as they cannot follow the evolution of the tidal debris and are restricted to cut-off criteria in energy or angular momentum space. Moreover, limits in computational power prevented us from carrying out such investigations by means of collisional $N$-body codes. But the recent improvements in computational speed of $N$-body codes and the availability of accelerator hardware (GRAPEs, GPUs) now allows us to study the dynamical evolution of star clusters with masses of up to several times $10^4\\msun$.\\\\  Thus, by computing various star clusters in a range of tidal conditions over several Gyr we investigate the structure of star clusters, and in particular the evolution of the region around the tidal radius - the transition region between the star cluster and the tidal tails - in terms of velocity dispersion and surface density profiles. The paper is organized as follows: first, a brief introduction to the topic of potential escapers is given in Sec.~\\ref{Sec:Pesc} as these stars play a key role in this investigation. Then a few methodological remarks will be made in Sec.~\\ref{Sec:Method} before we  come to the velocity dispersion profiles in Sec.~\\ref{Sec:VDP} and then to the surface density profiles in Sec.~\\ref{Sec:SDP}. In the last section we will give a summary with a brief discussion. ",
        "conclusions": "\\label{Sec:Conclusions} This systematic investigation of velocity dispersion profiles and surface density profiles of star clusters with masses of about $10^4 \\msun$ has shown how complex and versatile these profiles can be, depending on the circumstances of the clusters.\\\\  Furthermore we have shown the importance of potential escapers for these two quantities. Both profiles are influenced by this population of energetically unbound stars for radii larger than about 50\\% of the Jacobi radius and are completely dominated by it for radii larger than about 70\\% of the Jacobi radius (Fig.~\\ref{vdps} \\& \\ref{vdpse075}). \\citet{Baumgardt01} found that for clusters in a constant tidal field the fraction of potential escapers varies with the number of stars within a cluster approximately as $N^{-1/4}$. Thus, even if the cluster mass increases by a factor of 10, the fraction (and influence) of potential escapers will only be reduced by less than 50\\% compared to this investigation.\\\\  As a consequence of the potential escaper population, the velocity dispersion profiles in our sample do not show a clear separation between the bulk of the cluster and the tidal debris at the Jacobi radius, but rather show a smooth transition from cluster to debris. For mass estimates which are based on velocity dispersion measurements we therefore recommend to use stars within 50\\% of the Jacobi radius to minimise the effect of potential escapers.\\\\  Moreover, we argue that investigations on velocity dispersion profiles like \\citet{Drukier98}, \\citet{Scarpa03} or \\citep{Scarpa07} have to be interpreted with caution. Also, detecting a possible dark matter halo around globular clusters is less feasible due to the population of potential escapers. Based on our investigation we suggest that no deviation from Newtonian gravity in a star cluster has been detected so far, neither has there been evidence for a dark matter halo.\\\\  From our large set of surface density profiles we saw that in most cases the structure of extended star clusters can be split into core, bulk and tidal debris. \\begin{enumerate} \\item The core is the innermost part of the cluster which extends out to the core radius.  In most cases it is flat, and only for clusters with strong mass segregation or with an IMBH is it cuspy. Very concentrated clusters with $R_{hp}/R_J < 0.1$ have such a small core that fitting such clusters leads to problems with templates as well as with physical models (see also \\citealt{Baumgardt10}).\\\\\\\\ \\item The bulk contains most stars of the cluster and extends from the core to the tidal debris. We found that the bulk can be well represented by a King-like template. We named the radius at which such a template approaches zero density the edge of the bulk and not the tidal radius of the cluster, since we found it not to be consistent with the Jacobi radius. In the case of concentrated clusters, the edge radius is not even determined by tidal forces  at all (Tab.~\\ref{table5}).\\\\  Moreover, for clusters on eccentric orbits we found that the edge radius of a cluster adjusts to the mean Jacobi radius (Fig.~\\ref{rt8500ecc}) and not the perigalactic Jacobi radius, which has been assumed in earlier investigations (e.g. \\citealt{Innanen83}, \\citealt{Fall01}), but has never been checked by self-consistent calculations on a star-by-star basis. Since nearly all globular clusters of the Milky Way are more compact and massive than our test cluster, and thus should be less influenced by tidal variations, this finding should also hold for most of these clusters.\\\\  The edge radius of extended and moderately concentrated clusters, as fitted by the King template, lies within about 10\\% accuracy at ~90\\% of the mean Jacobi radius (Tab.~\\ref{table3}-\\ref{table4c}). But we found the King template to be a bad representation of the true profile in the case of non-circular cluster orbits because it is significantly influenced by the tidal debris and by the potential escaper population. Furthermore the accuracy of the results strongly depends on the radial coverage of the cluster. We therefore created an enhanced version of the King template which has more flexibility in the core and has an additional tidal debris term, given by a power-law. With this KKBH template we achieved a higher stability over all fit ranges. Furthermore, we were able to measure the edge of the bulk more accurately and found it to lie at 80\\% of the Jacobi radius. This discrepancy is easily understandable if we take into account the fact that the edge radius of a template fits the azimuthally averaged mean of a cluster's tidal surface, whereas the Jacobi radius is by definition the semi-major axis of the equipotential surface.\\\\  Furthermore we found that the edge radii from concentrated clusters with $R_{hp}/R_J < 0.1$ cannot be extracted properly with currently available templates. The higher the concentration of the cluster the smaller is the ratio of the fitted edge radius to the theoretically evaluated Jacobi radius (Tab.~\\ref{table5}).\\\\\\\\  \\item The tidal debris falls off like a power-law with a slope of about -4 to -5, rather than the  expected dependence of  $R^{-1}$. This is due to the epicyclic motion of the stars in the tidal tails, which is most pronounced in the vicinity of the cluster.\\\\  For clusters on eccentric orbits, at a time shortly before reaching apogalacticon, the slope was found to deviate for a short time to a value of about -2 as a consequence of orbital compression of the tails (Fig.~\\ref{eta8500ecc} \\& \\ref{sdp17000}). We suggest that this enhanced slope can be used to identify star clusters which are close to or at the apogalacticon of their orbit. The most prominent example of such a cluster is the Milky Way globular cluster Palomar 5 which is well known to be close to the apogalacticon of its orbit \\citep{Dehnen04}, and which shows a power-law slope of -1.5 outside its assumed Jacobi radius \\citep{Odenkirchen03}. Moreover, the MW globular clusters Palomar 13, AM 4 and Whiting 1 may be further candidates for being close to apogalacticon, since their surface density profiles show power-law slopes of about -1.8 outside the Jacobi radius, which has been attributed to heavy ongoing mass loss by the authors of the corresponding investigations \\citep{Cote02, Carraro07, Carraro09}.\\\\  Furthermore we showed that due to the population of potential escapers the power-law slope of the tidal debris  begins even before  the Jacobi radius is reached. Thus, the bulk and the tidal debris partially overlap. With our KKBH template we found the debris to dominate for radii larger than about 50\\% of the Jacobi radius (Tab.~\\ref{table3}-\\ref{table4c}). \\end{enumerate}  With the help of our comprehensive set of clusters it is possible to tailor a star cluster template to the  facts we have found. Our KKBH template seems to be a good first step but, like all other available templates, shows discrepancies for concentrated clusters. Of course, templates are less attractive than physically motivated models. None of the existing models, like \\citet{King66} and \\citet{Wilson75}, however, account for the above mentioned facts.  We suggest that a general, physically motivated model for star clusters should include a term for mass segregation and should include a potential escaper population when being fitted to observations. Finally, the possible influence of (compressed) tidal tails on the surface density profiles of star clusters should be kept in mind when fitting surface density profiles.\\\\  Moreover, the fact that all clusters in the sample showed a constant edge of the bulk with time (Fig.~\\ref{rt8500ecc}) and that stars which get unbound escape from a cluster with a delay, supports our theoretical treatment of epicyclic overdensities in tidal tails for clusters in time-dependent tidal fields in \\citet{Kuepper09}."
    },
    "1005/1005.3057_arXiv.txt": {
        "abstract": "{We present a resolved dust analysis of three of the largest angular size spiral galaxies, NGC 4501 and NGC 4567/8, in the \\textit{Herschel} Virgo Cluster Survey (HeViCS) Science Demonstration field. \\textit{Herschel} has unprecedented spatial resolution at far-infrared wavelengths and with the PACS and SPIRE instruments samples both sides of the peak in the far infrared spectral energy distribution (SED). We present maps of dust temperature, dust mass, and gas-to-dust ratio, produced by fitting modified black bodies to the SED for each pixel. We find that the distribution of dust temperature in both systems is in the range $\\sim$19 - 22~K and peaks away from the centres of the galaxies. The distribution of dust mass in both systems is symmetrical and exhibits a single peak coincident with the galaxy centres. This Letter provides a first insight into the future analysis possible with a large sample of resolved galaxies to be observed by \\textit{Herschel}.} ",
        "introduction": "Infrared data have been widely used to determine the composition and distribution of dust in galaxies since the launch of the Infrared Astronomical Satellite (IRAS) in the 1980s. However, dust masses for nearby galaxies calculated from IRAS 60 and 100\\micron\\ measurements were found to be a factor of ten lower than expected when compared to the Milky Way gas-to-dust ratio of 100-200 \\citep{devereux1990}. The Milky Way dust mass was calculated by measuring the depletion of metals from the gaseous phase of the interstellar medium (ISM) and by comparing gas column densities to dust extinction \\citep{whittet03}, its value implying that most of the dust-mass emits radiation at wavelengths longer than 100\\micron\\ \\citep[e.g.,][]{devereux1990}. Analyses of data from the {\\em Spitzer Space Telescope\\/} have determined gas-to-dust ratios of $\\sim$150 \\citep[e.g.,][]{draine07}. However, these analyses were generally limited by the number, and low signal-to-noise ratio of data at wavelengths greater than 160\\micron. These studies therefore had difficulty detecting emission from dust with temperatures lower than 15~K, and their results were biased towards warmer dust temperatures and lower masses. Owing to the difficulties at these wavelengths, high resolution studies of galaxies have been previously limited. These studies are important for understanding how dust interacts with the other phases of the ISM, the sources of dust heating and how the distribution and temperature of dust varies with morphology.  The \\textit{Herschel Space Telescope} \\citep{pilbratt10}, with its two photometric instruments, constrains both sides of the peak in the far infrared spectral energy distribution (SED, see Fig. \\ref{fig:SEDs}). The Photodetector Array Camera \\& Spectrometer (PACS, \\citealp{Poglitsch10}) has three photometric bands at 70, 100, and 160\\micron\\ at superior angular resolutions to those provided by \\Spitzer. The Spectral and Photometric Imaging Receiver (SPIRE, \\citealp{Griffin10}) also has three photometric bands observing simultaneously at 250, 350, and 500\\micron\\ with high sensitivity and angular resolution.  The Herschel Virgo Cluster Survey (HeViCS\\footnote{More details on HeViCS can be found at http://www.hevics.org}) is a \\textit{Herschel} Open Time Key Project that will observe $\\sim$64~deg$^2$ of the Virgo cluster. This will provide a large sample of resolved galaxies because about 48 late-type galaxies will be observed by HeViCS with optical diameters larger than 3\\arcmin. In this Letter, we present an insight into what will be possible with the full HeViCS survey by applying a resolved dust analysis to infer dust temperatures, surface densities, and gas-to-dust ratios for NGC 4501 and NGC 4567/8. These galaxies were chosen because they are among the largest angular size systems in the HeViCS Science Demonstration Phase (SDP) field. In Sect. 2, we present the observations and data reduction, and in Sects. 3 and 4, we present our analysis and results, respectively. ",
        "conclusions": "We have fitted SEDs on a pixel-by-pixel basis to investigate the dust mass, temperature, and gas-to-dust ratios of NGC 4501 and NGC 4567/8. We have measured dust surface densities that peak at the centre of the galaxy and decrease towards the outer regions. In contrast, the temperature distribution is asymmetric, higher temperatures peaking at some distance away from the centre of the galaxy and then decreasing with increasing radius towards the outskirts. Once the full HeViCS observations are complete, we will be able to extend the analysis to larger radii and lower dust surface densities. For $\\sim$48 late-type galaxies with optical radii greater than 3\\arcmin\\ to be observed with HeViCS, we will be able to begin to tackle questions about the dominant source of dust heating, how the dust interacts with the other phases of the ISM, and how morphology influences the dust distribution and temperature. By combining the results from HeViCS and the Herschel Reference Survey \\citep{boselli10}, we will be able to study how environment affects the dust distribution and other properties of a galaxy."
    },
    "1005/1005.5098_arXiv.txt": {
        "abstract": "{The recently discovered exoplanet Gl581d is extremely close to the outer edge of its system's habitable zone, which has led to much speculation on its possible climate. We have performed a range of simulations  to assess whether, given  simple  combinations of chemically stable greenhouse gases, the planet could sustain liquid water on its surface. For best estimates of the surface gravity, surface albedo and cloud coverage, we find that less than 10 bars of CO$_2$ is sufficient to maintain a global mean temperature above the melting point of water. Furthermore, even with the most conservative choices of these parameters, we calculate  temperatures  above the water melting point  for CO$_2$ partial pressures greater than about 40 bar. However, we note that as Gl581d is probably in a tidally resonant orbit, further simulations in 3D are required to test whether such atmospheric conditions are stable against the collapse of CO$_2$ on the surface.} ",
        "introduction": "In 2007, radial velocity measurements were used to discover two new planets in the Gl581 system \\citep{udry2007}. These planets have captured much attention both in the community and among the general public, as their minimum masses were measured to be below 10 $M_{Earth}$, and they are close to the edges of their system's nominal `habitable zone', i.e., the loosely defined orbital region in which planets can sustain liquid water on their surfaces. The first planet, Gl581c, which is closer to its star and was the first discovered, was initially estimated to be potentially habitable based on its equilibrium temperature $T_{eq}=320$ K, using an Earth-like planetary albedo 0.29.  In contrast, the second planet Gl581d  has an equilibrium temperature $T_{eq}= 195$ K for  an albedo of 0.2, which suggests it may be too cold to sustain surface liquid water. However, these analyses neglect any possible warming of the surface due to the planet's atmosphere.  In the first detailed assessment of the potential habitability of these planets, \\citet{selsis2007} reviewed a variety of factors that could influence their climates. They concluded that based on standard assumptions of atmospheric warming by a mixture of CO$_2$ and H$_2$O (with possible regulation of CO$_2$ via the carbonate-silicate cycle), Gl581c was unlikely to be habitable, while for Gl581d,  the situation was much less clear. According to \\citet{kasting1993}, the outer edge of the habitability zone is most likely the distance at which CO$_2$ condensation begins to occur on the surface of the planet. However, CO$_2$ condensation in the atmosphere leads to the formation of CO$_2$ clouds, which can cause a strong warming effect due to the scattering of infrared radiation \\citep{forget1997}. Hence it was concluded that further climate simulations were required.  To investigate the possible climate of Gl581d under a range of conditions, we have performed one-dimensional radiative-convective  calculations. In Section \\ref{sec:method} we discuss the model we used, while in Section \\ref{sec:results} we present our results for varying atmospheric compositions, surface albedos, gravity and cloud coverage. We also present some simple three-dimensional simulations that highlight the limitations of the one-dimensional globally averaged approach. In Section \\ref{sec:discuss}, we discuss the implications of our results and suggest directions for future research. ",
        "conclusions": "\\label{sec:discuss}  Through one-dimensional radiative-convective climate modeling, we have found that as little as 5 bars of CO$_2$ ($\\sim$2.5 bar equivalent column amount under Earth gravity) may be sufficient to maintain a global mean temperature above the melting point of water. For CO$_2$ partial pressures greater than 40 bar ($\\sim$13 bar equivalent column amount under Earth gravity), we found that the mean surface temperature of Gliese 581d is high enough to allow surface liquid water even when the most conservative values of gravity, surface albedo and N$_2$ partial pressure are chosen, although  strong dependence on the total H$_2$O cloud opacity was observed.  We tested the dependence of surface temperature on a wide range of atmospheric and geophysical parameters. Because the M-class spectrum of Gliese 581 is red-shifted with respect to that of the Sun, heating due to near-IR absorption in the upper atmosphere increases in our simulations, while Rayleigh scattering is much less important. This is in agreement with the general habitability study of \\citet{kasting1993}, but we find the difference to be even stronger than was reported there, due to the low effective temperature of Gl581. We also estimate an modest warming effect due to CO$_2$ clouds (up to 30 K at $p_{CO_2}$ = 50 bar), dependent on details of the cloud microphysics, although the efficiency of this process is reduced compared to the case for G-class stars.  We find that increasing gravity primarily decreases the CO$_2$ column amount at which the maximum surface temperature occurs. Unsurprisingly, increased surface albedo decreases the mean surface temperature, although if there is cloud coverage, this effect is relatively small for the high CO$_2$ pressures at which the surface starts to become habitable. Increasing the partial pressure of N$_2$ has little effect below partial pressures $p_{N_2}$ of tens of bars. Finally, water vapour increases the surface temperature by a moderate amount below $T_s = 273 $ K, after which its role becomes increasingly important. We briefly assessed the impact of H$_2$O clouds, and found that they warmed (cooled) the surface at low (high) pressures, due to the interplay between their effects on incoming stellar and outgoing thermal radiation.  As discussed by \\citet{selsis2007}, N$_2$-CO$_2$-H$_2$O atmospheres likely only represent a fraction of the possible range for Earth-like / super-Earth extrasolar planets.  While the presence of radiatively inactive gases such as argon would be of little importance to the calculations discussed here, small amounts of other greenhouse gases such as methane would result in increased absorption, if such gases could remain chemically stable in the atmosphere. Unless the near-IR absorption of these gases were large enough to cause a strong temperature inversion in the upper atmosphere, this would further increase the surface temperature of the planet. The presence of other clouds or haze layers could alter the radiative balance in less predictable ways, depending on their scattering / absorption properties and height in the atmosphere. Unfortunately, effects of this kind are difficult to constrain further until direct atmospheric observations become available.  The results discussed here clearly have interesting implications, as they show that Gl581d could be the first discovered habitable exoplanet. They are also testable by future observations - dense CO$_2$ atmospheres have recognizable spectral signatures that could be detected, for example, using the proposed Darwin or TPF missions. However, there is an important limitation to this study: all simulations performed were one-dimensional. As we noted in Section \\ref{subsec:3Dnoatm}, calculations of the \\emph{globally averaged} surface temperature neglect variations due to local changes in stellar insolation (as well as topography and other effects). This has serious implications for dense, relatively cold CO$_2$ atmospheres, as if the dark side temperature of the planet is too low, CO$_2$ will condense on the surface. To test whether a dense atmosphere on Gl518d would be stable against the collapse of CO$_2$ in this way, we plan to repeat the calculations reported here in the future using a three-dimensional climate model."
    },
    "1005/1005.2392_arXiv.txt": {
        "abstract": "Inertial waves governed by Coriolis forces may play an important role in several astrophysical settings, such as eg.  tidal interactions, which  may occur in extrasolar planetary systems and close binary systems, or in rotating compact objects emitting gravitational waves. Additionally, they are of interest in other research fields, eg. in geophysics.  However, their analysis is complicated by the fact that in the inviscid case the normal mode spectrum is either everywhere dense or continuous in any frequency interval contained within the inertial range. Moreover, the equations governing the corresponding eigenproblem are, in general, non-separable.  In this paper we develop a consistent WKBJ formalism, together with a formal first order perturbation  theory for calculating the properties of the normal modes  of a uniformly rotating coreless body (modelled as a polytrope and referred hereafter to as a planet) under the assumption of a spherically symmetric structure. The eigenfrequencies, spatial form of the associated eigenfunctions  and  other properties we  obtained analytically using the WKBJ eigenfunctions are in good agreement with corresponding results obtained by numerical means  for a variety of planet models even for global modes with a large scale distribution of perturbed quantities. This indicates that even though they are embedded in a dense spectrum,  such modes can be identified and followed as model parameters changed and that first order perturbation theory can be applied.  This is used to estimate corrections to the eigenfrequencies as a consequence of the anelastic approximation,  which we argue here to be small when the rotation frequency is small. These are  compared with simulation results in an accompanying paper with a good agreement between theoretical and numerical results.  The results reported here may provide a basis of theoretical investigations of inertial waves in many astrophysical and other applications, where a rotating body can be modelled as a uniformly rotating barotropic object, for which the density has, close to its surface, an approximately power law dependence on distance from the surface. ",
        "introduction": "In astrophysical applications inertial waves that can exist in rotating bodies may be excited by several different physical mechanisms, most notably through tidal perturbation by a companion (eg. Papaloizou \\& Pringle 1981, hereafter PP) or in the case of compact objects through secular instability arising through gravitational wave losses (eg. Chandrasekhar 1970, Friedman \\& Schutz 1978, Andersson 1998, Friedman $\\&$ Morsink 1998). They also can play a role in other physical systems. For example, they can also be excited by several mechanisms in the Earth's  fluid core with possible detection being announced (Aldridge $\\&$ Lumb 1987).   For  rotating planets and stars  that have  a barotropic equation of state these  wave modes are governed by Coriolis forces and so have oscillation periods that are comparable to the rotation period. They are accordingly readily excited by tidal interaction with a perturbing body when the characteristic time associated with the orbit is comparable to the rotation period, which is expected naturally when the rotation period and orbit become tidally coupled. They may then play an important role in governing the secular orbital evolution of the system.  Inertial modes excited in close binary systems in circular orbit were considered by PP and Savonije \\& Papaloizou  (1997). Wu (2005)a,b considered the excitation of inertial modes in Jupiter as a result of tidal interaction with a satellite and excitation as a result of  a parabolic  encounter of a planet or star with a central star was studied by Papaloizou \\& Ivanov (2005), hereafter referred to as PI and Ivanov \\& Papaloizou (2007), hereafter referred to as  IP. The latter work was applied to the problem of circularisation of extrasolar giant planets starting with high eccentricity. In that  work the planet  was  assumed coreless. Ogilvie \\& Lin  (2004) and Ogilvie (2009) have considered the case of a cored planet in circular orbit around a central star and found that inertial waves play an important role.  The importance of the role played by  inertial waves in the  transfer of the rotational energy of a rotating neutron star to gravitational waves via the Chandrasekhar-Friedman-Schutz  (CFS) instability was pointed out by Andersson (1998). Later studies mainly concentrated on physical mechanisms of dissipation of energy stored in these modes that limit amplitudes of the modes, and, consequently, the strength of the  gravitational wave signal. In these studies either numerical methods or simple local estimates of properties of inertial modes were mainly used, see eg. Kokkotas (2008) for a recent review and references.   An analytical treatment of problems related to inertial waves, such as eg. finding normal mode spectra and  eigenfunctions, and coupling them  to other physical fields, etc., is difficult due to a number of  principal complicating technical issues.   In particular, the dynamical equations governing the  perturbations of a rotating body (called planet later on) are, in general, non-separable, for compressible fluids. When such fluids are considered and rotation is assumed to be small,  a low frequency anelastic approximation that filters out the high frequency modes is often used (see eg. PP). This simplifies the problem to  finding solutions to leading order in the  small parameter $R_{*}^{3}\\Omega^{2}/(GM)$, where $\\Omega $ is the rotation frequency, $G$ is the constant of gravity and $M_{*}$, $R_{*}$ are the mass and radius of the planet. In this approximation eigenfrequencies of inertial modes are proportional to $\\Omega ,$ while the form  of the spatial distribution  of  perturbed quantities does not depend on the rotation rate. However, even when this approximation is adopted, the problem is, in general, non-separable apart from models with a  special form of  density  distribution, see Arras et al (2003), Wu (2005)a and below.   Additionally, the problem of   calculating the inertial mode spectrum and its response to tidal forcing is complicated by the fact that in the inviscid case  the spectrum is either everywhere dense or continuous in any frequency interval it spans (Papaloizou \\& Pringle, 1982). This is in contrast  to the situation of, for example, high frequency $p$ modes, which are discrete with well separated eigenvalues. When the anelastic approximation  is adopted the singular ill posed nature of the inviscid eigenvalue problem is seen to come from the fact that the  governing equation is hyperbolic and the nature of the spectrum is  determined by the properties of the characteristics (eg. Wood 1977).  A discrete spectrum is believed to occur when there are no such trajectories that define  periodic attractors. Otherwise the  inviscid spectrum is continuous. Then, when a small viscosity is introduced the spectrum becomes  discrete but normal modes have energy focused onto wave attractors (see eg. Ogilvie \\& Lin 2004). Given these complexities it is desirable to work with and compare a variety of analytical and numerical approaches.  Coreless  inviscid rotating planets with an assumed spherical or ellipsoidal shape have a discrete but everywhere dense spectrum that makes difficulties for example with mode identification and application of  standard perturbation theory. However, numerical work indicates that there are well defined global modes that can be identified and followed through a sequence of models (eg.  Lockitch \\& Friedman, 1999, hereafter  LF, and   PI). In this paper we investigate the inertial mode spectrum of a uniformly rotating  coreless barotropic  planet or star and its tidal response by a WKBJ approach coupled with first order perturbation theory   and compare its eigenvalue predictions with numerical results obtained by a variety of authors  and find good agreement  apart from some unidentified WKBJ modes that are near the limits of the spectrum and for which the perturbation theory appears not to work. For the identified modes we also find remarkably  good agreement for the form of the eigenfunctions. This indicates that they can be represented at  low resolution with  small scale phenomena being unimportant,  meaningful mode identification (in that the modes can be followed from one model to another) and at least first order perturbation theory works for these modes.  This is also confirmed in a following paper  (hereafter referred to as PIN) where we investigate the inertial mode spectrum and its tidal response by numerical solution of an initial value problem {\\it  without the anelastic approximation}. We are able to confirm the validity of the anelastic approximation and the applicability of the first order perturbation theory developed here  for demonstrating this as well as estimating eigenvalues. Thus a suggestion  of Goodman \\& Lackner (2009) that tidal interaction might be seriously overestimated by use of the anelastic approximation is not confirmed.  A WKBJ approach to the same problem was also considered by  Arras et al (2003) and Wu (2005)a. However, in this work  only terms of leading order in an expansion in inverse powers of a large WKBJ parameter $\\lambda $ (see the text below for its definition) were taken into account and treatment of perturbations near the surface  and close to the rotational axis were oversimplified. As a consequence, although their results  are correct in the formal limit $\\lambda \\rightarrow \\infty, $ they cannot be used to make a correspondence between WKBJ modes  and those  obtained  numerically, or an approximate description of modes with a scale that is not very small. In this paper we treat the problem in a more  extended way,  considering  terms of the next $O(\\lambda^{-1})$ order together with an accurate treatment of perturbations near the surface and close to the rotation axis. Additionally, we consider a frequency correction of the next order, $O(\\lambda^{-2}),$  for modes having non-zero azimuthal number, $m$.  We checked results obtained with use of the WKBJ formalism against practically all numerical data existing in the literature finding good agreement in practically all cases. Therefore, we can assume that our formalism may be applied to provide an approximate analytic description of inertial modes,  including those with large scale variations, where the WKBJ approach might be expected to be  invalid. Also, different quantities associated with the modes may be described within the framework of our formalism or its natural extension, such as the tidal overlap integrals (see PI and IP), quantities determining the growth rate due to the CFS instability and decay of inertial waves due different processes, eg. by non-linear mode-mode interactions (see eg. Schenk et al 2002, Arras et al 2003). Thus, the formalism developed here may provide a basis for the  analytic  treatment of inertial waves in many different astrophysical applications.   The plan of the paper is as follows. In section  \\ref{sec2} we briefly review the basic equations and their linearised form for a uniformly rotating barotropic planet or star. In section \\ref{anelastic} we go on to consider these in the anelastic approximation which is appropriate when the rotation frequency of the star is very much less than the critical or break up rotation frequency. We give a simple physical argument why we expect this approximation to be valid in this limit even when the sound speed tends to a small value or possibly zero at the surface of the configuration. In section \\ref{sec2.5} we give a brief discussion about when discrete normal modes may be expected to occur such as in the case of a coreless slowly rotating planet with surface boundary assumed to be either spherical  or ellipsoidal. We then  present a formal first order perturbation theory that can be used to estimate corrections to eigenfrequencies occurring as either a consequence of  terms neglected in the WKBJ approximation or the anelastic approximation. The latter application is tested by a direct comparison with the results of numerical simulations in PIN. Section \\ref{sec2.6} concludes with a brief account of the form of the anelastic equations in pseudo-spheroidal coordinates in which they become separable for density profiles of the form $\\rho \\propto (1-r^2/R_*^2)^{\\beta},$ where $r$ is the local radius, $R_*$ is the surface radius and $\\beta$ is a constant. (Arras et al 2003, Wu 2005a).   In section \\ref{sec3} we develop a WKBJ approximation for calculating the normal modes which is based on the idea that in the short wavelength limit these modes coincide with those appropriate to separable cases which include the  homogeneous incompressible sphere as a well known example. Solutions of a general WKBJ form appropriate to the interior of the sphere are matched to solutions appropriate to the surface regions where they become separable which is the case when the density vanishes as a power of the distance to the boundary as is expected for a polytropic equation of state. This matching results in an  expression for the eigenfrequencies given in section \\ref{sec3.5}.  In section \\ref{surface} we go on to develop expressions for the eigenfunctions appropriate to  any location in the planet including the rotation axis and the critical latitude region where one of the inertial mode characteristics is tangential to the planet surface. These solutions are then used to obtain corrections to the eigenfrequencies resulting from density gradient terms neglected in the initial WKBJ  approximation in section \\ref{sec3.9}. In section \\ref{sec4} we compare the corrected  eigenfrequencies obtained from the WKBJ approximation with those obtained numerically by several different authors who used differing numerical approaches and find good agreement even for global modes. A similar comparison with the results of  numerical simulations for a polytropic model with  positive results  is reported in PIN. We also  compare the forms of the eigenfunctions with those obtained in Ivanov \\& Papaloizou (2007) and find a  good agreement even for global modes.  {Finally in section \\ref{sec5.1} we discuss our results in the context of the evaluation of the overlap integrals that occur in evaluating the response to tidal forcing. We show that these vanish smoothly in the limit  that the polytropic index tends to zero and we indicate that they vanish at the lowest WKBJ order and  are  thus expected to vanish rapidly as the order of the mode increases. We go on to summarize our conclusions in section \\ref{Conclu}.} ",
        "conclusions": "\\label{sec5}  \\subsection{Overlap integrals}\\label{sec5.1}  { As we pointed out in the Introduction, integrals of the form \\begin{equation} \\hat Q_{k}={Q_{k}\\over \\sqrt N_{k}} \\quad {\\rm with}\\quad Q_{k}= ({\\rho \\over c_{s}^{2}}W_{k}|\\Phi), \\label{eqn84} \\end{equation} where $W_{k}$ corresponds to a particular eigenmode and $\\Phi $ is some smooth function, appear in astrophysical applications of the theory developed in this paper. In particular, as was discussed in PI and IP, integrals of this type enter in expressions for the transfer of energy and angular momentum transferred during  the periastron passage of a massive perturber. These apply to the case when the spectrum of normal  modes is discrete and they  involve integrals of form (\\ref{eqn84}), where $\\Phi=P_{2}^{m}r^{2}$  with $P_{2}^{m}$   being the associated Legendre function. Assuming that  $W_{k}$ varies on a small spatial scale while the function $\\Phi $ is smoothly varying such integrals may, in principal, be evaluated using our formalism  with help of a theory of asymptotic evaluation of multidimensional integrals, see eg.  Fedoryuk (1987), Wong (1989).  However, some important integrals of form (\\ref{eqn84}) require an extension of our formalism, which can provide a smooth matching of the solution close to the surface to the WKBJ solution in the inner part of the planet that is valid  at  the next orders in inverse powers of $\\lambda $. This is due to cancellations of leading terms in corresponding asymptotic series. Since this problem appears to be a rather generic one we would like to discuss it here in  more detail  for the important case when $\\Phi =P^{2}_{2}r^2=3\\varpi^{2}$. The overlap integral of this type determines excitation of the $m=2$ modes  which are the most important for the tidal problem (eg. PI, IP and see also PIN). Explicitly, we have in this case \\begin{equation} Q_{k}= 3\\int dV \\left(\\varpi^2 {\\rho \\over c_{s}^{2}}W_{k}\\right), \\label{eqn84a} \\end{equation} where $dV=dz\\varpi d\\varpi $. Note that this integral must converge to zero in the incompressible limit $n\\rightarrow 0$ as in this case it is well known that inertial modes are not excited in the anelastic approximation. This fact, however, is not obvious for the integral written in the form (\\ref{eqn84a}) because  close to the surface we have \\begin{equation} {\\rho \\over c_{s}^{2}}\\approx nCx^{(n-1)}, \\label{eqn85} \\end{equation} with the constant $C$ converging to a nonzero value  as $n \\rightarrow 0.$  Therefore, as the eigenfunctions are regular, the integral $Q_{k}/n$ has a logarithmic divergence at the surface of the planet as $n \\rightarrow 0.$  This raises the possibility that the overlap integral might converge to a nonzero value or behave pathologically as the incompressible limit is approached.  In order to show that, in fact, this is not so and $Q_{k}(n\\rightarrow 0)\\rightarrow 0$ in a smooth manner, let us consider some fiducial models having the property that the quantity \\begin{equation} \\omega_0^2 = -\\frac{c^2_s}{r \\rho }\\frac{d\\rho }{dr}\\label{Good1} \\end{equation} is constant.  For models in  hydrostatic equilibrium  under their own gravity, constancy of $\\omega_0^2$ implies that ratio $M(r)/r^{3}$, where $M(r)$ is the mass interior to the radius $r,$  is constant. The model must accordingly be  incompressible. Goodman $\\&$ Lackner (2009) obtained a wider class of models in hydrostatic equilibrium   under a fixed quadratic gravitational potential. Because the potential is fixed independently of the mass distribution and there are no constraints on the equation of state, such models may be constructed for an arbitrary density distribution.  Now let use consider the  integral \\begin{equation} Q_{fud}= 3\\int dV \\left (\\varpi^2 {\\rho \\over c_{s}^{2}}\\omega_0^2W_{k}\\right). \\label{eqn84b} \\end{equation} For the fiducial models described above this is identical to the overlap integral (\\ref{eqn84a}) where we note that  we may adopt natural units  such that the constant $\\omega_0^2, $ which should be identified with the surface value of $GM(r)/r^3$ is equal to unity in that case. More generally the integrand in (\\ref{eqn84b}) can be transformed using equation of hydrostatic equilibrium (\\ref{Good1}) such that \\begin{equation} Q_{fud}= -3\\int dV \\left ({\\varpi^2\\over r} {d\\rho \\over dr}W_{k}\\right)=-3\\int dV \\left (\\varpi {\\partial \\rho \\over \\partial \\varpi}W_{k}\\right). \\label{eqn84c} \\end{equation}  Now let us consider equation (\\ref{eq 41}) for free normal modes in the anelastic approximation  by setting $\\sigma =\\sigma_k$ and the right hand side of this equation to zero. Then we multiply it by $\\varpi^2$, set $m=2$, and integrate over $dV$. After removing derivatives of $W_k$ by integrating by parts, assuming that the density vanishes at the surface boundary, it is easy to see that it follows from (\\ref{eq 41}) that $Q_{fud}=0$ in the anelastic approximation provided $\\sigma_k \\ne 2\\Omega.$ This means, in particular, that $m=2$ inertial waves cannot be excited in the  Goodman $\\&$ Lackner (2009) models in this approximation as was found by these authors  when compressibility was fully taken into account (see  PIN for additional discussion).  Using the fact that $Q_{fud}=0$ we may rewrite (\\ref{eqn84a}) for models under their own self-gravity quite generally, adopting natural units,  as \\begin{equation} Q_{k}=Q_{k}-Q_{fud}= 3\\int dV \\varpi^2 {\\rho \\over c_{s}^{2}}\\left(1-\\frac{M(r)}{r^3}\\right)W_{k}. \\label{eqn84d} \\end{equation} Taking into account that the factor in the brackets is proportional to $x$ for small $x$, and, accordingly, in this limit the integrand is proportional to $nx^n$ we see that now the logarithmic divergence of $Q_k/n$ disappears and, therefore, it is clear from the representation (\\ref{eqn84d}) that the overlap integral indeed  smoothly tends  to zero in the limit $n\\rightarrow 0$.  The theory of asymptotic evaluation of integrals of the form (\\ref{eqn84d}) tells that the values of such integrals are determined either by inner stationary points, where gradient of the WKBJ phase vanishes or contributions close to the surface or other parts of the integration domain, where the WKBJ approximation is not valid. From the expression of the $W_{k}$ in the WKBJ regime (\\ref{eqn30}) it follows that there are no stationary points in the inner region of the planet. Considering the regions close to the surface it appears to be reasonable to assume that the leading contribution is determined by the region close to the critical latitude, where a 'hot spot' is observed in distributions of $W_k$, see the previous section. In this region the quantities $\\delta_1=x_1-\\mu$ and $\\delta_2=\\mu-x_2$ are small. We can use them as new integration variables in (\\ref{eqn84d}) with help of (\\ref{eqn66}), separate the contribution of this region to the integral by  introduction of the functions $\\eta[\\delta_{1,2}]$ defined in section 3.7 in the integrand, and decompose the quantities in front of $W_k$ in powers of $\\delta_{1,2}$ taking into account that $x\\propto \\delta_{1}\\delta_2$ and $\\varpi^{2}\\approx (1-\\mu^{2})$ in the leading order. Assuming that $W_{k}\\approx \\bar W(\\delta_{1}) \\bar W(\\delta_{2})$, where $\\bar W(y)$ is given by equation (\\ref{eqn60}), it is easy to see that the leading contribution to (\\ref{eqn84}) is given by a symmetric combination of two integrals involving Bessel functions \\begin{equation} Q_{k}\\propto I_{1}(\\delta_{1})I_{2}(\\delta_{2})+ I_{2}(\\delta_{1})I_{1}(\\delta_{2}), \\label{eqn84g} \\end{equation} where \\begin{equation} I_{1}(y)=\\int_{y=0}dy \\eta[y]y^{n}J_{(n-1)/2}(\\kappa y), \\quad I_{2}(y)=\\int_{y=0}dy \\eta[y]y^{(n+1)}J_{(n-1)/2}(\\kappa y), \\label{eqn84j} \\end{equation} where it is assumed that $\\eta[y >y_{*}]=0$, $y_{*}$ lies within the range of integration and $1/\\kappa \\ll y_{*} \\ll 1$. As was shown by Larichev (1973) the integral $I_{1}(y)=0$ for any particular form of the function $\\eta[y]$. Thus, the leading order contribution to the overlap integral from the surface region close to the critical latitude is equal to zero. In principal, one can look for the next order terms. However, in this case our simple approach to the problem seems to be inadequate since eg. the assumption that $W_{k}$ can be represented as a product of two functions separately depending on the coordinates may be broken at this level, etc.. A more accurate approach is left for a future work. We note, however, that this cancellation means that the overlap integral should decay rapidly with increasing $\\lambda,$ possibly being  inversely proportional to a large power of $\\lambda.$ This may qualitatively explain why a small number of relatively large scale modes are significantly excited by dynamic tides, see PI, IP and PIN.}     \\subsection{Conclusions}\\label{Conclu}   In this paper we  have developed a WKBJ approximation, together with a formal first order perturbation approach for calculating the normal modes  of a uniformly rotating coreless planet under the assumption of a spherically symmetric structure. Matching of the general WKBJ form valid  in the interior to separable solutions valid near the surface resulted in expressions for eigenfunctions that were valid at any location within the planet together with an  expression for the associated eigenfrequencies given in section \\ref{sec3.5}.  Corrections as a result of density gradient terms neglected in the initial WKBJ approach were also obtained from formal first order perturbation theory.  The corrected WKBJ eigenfrequencies obtained using the WKBJ eigenfunctions were compared with results obtained numerically by several different authors and found to be in good agreement,  away from the limits of the inertial mode spectrum  where identifications could be made,  even for modes with a global structure. We also compared the spatial forms of the eigenfunctions with those obtained using the spectral method described in IP finding similar good agreement.  This is consistent with the idea that these global modes can be identified and that first order perturbation theory works even though they are embedded in a dense spectrum.  In further support of this, the formal first order perturbation theory developed here is subsequently  used to estimate corrections to the eigenfrequencies as a consequence of  the anelastic approximation and is then compared with simulation results for a polytropic model with $n=1$ in PIN. These different approaches are found to be in  agreement  for small enough rotation frequencies and also indicates  that, as implied by the simplified discussion in section \\ref{anelastic} of this paper, that  corrections as a result of the anelastic approximation are never very significant for the models adopted.  Our results show that the problem of finding eigenfrequencies and eigenvalues of inertial modes allows for an approximate analytical treatment, even in the case of modes having a large scale distribution of perturbed quantities.  Although we consider only the case of a polytropic planet, our formalism can be applied to a much wider context. Indeed, the approach developed here is mainly determined by the form  of the density close to planet's surface, where we assume that it is proportional to a power of distance from the surface. Thus, we expect that our main results remain unchanged for any density distribution, which is approximately power-law close to the surface. In particular, according to our results, all models of type having approximately the same behaviour of the density distribution close to the surface should have approximately the same eigenspectrum.  The formalism developed here can be extended for an approximate analytic evaluation of different quantities associated with inertial modes, such as overlap integrals characterising interaction of inertial waves with different physical fields, growth rates due to the CFS instability and decay rates due to various viscous interactions and non-linear mode-mode interaction. It may provide a basis for a perturbative analytic analysis of more complicated models, such as realistic models of star and planets flattened by rotation or models of relativistic stars.  As we discussed above, for a given value of WKBJ order, $l,$ some modes are identified with  modes obtained  numerically  while others remain unidentified. Eigenfrequencies of the unidentified modes are always either situated close to the boundaries of the frequency range allowed for inertial modes, $\\sigma =\\pm 2\\Omega $ or situated close to the origin $\\sigma=0$. We believe that our theory is not applicable to these modes, and they  develop  a small scale  contribution controlled by the closeness of the  positions of their eigenfrequencies to $\\sigma=0$ and $\\pm 2\\Omega, $ and thus effectively move to higher order than allowed for. Accordingly we de not consider these modes when comparing our results with results of direct numerical calculations of the excitation of inertial waves due to a tidal encounter reported in PIN."
    },
    "1005/1005.5267_arXiv.txt": {
        "abstract": "{A crucial issue in star formation is to understand the  physical mechanism by which mass is  accreted onto and  ejected  by a  young  star. To derive key  constraints on the launching point of the jets and on the geometry of the winds, the visible spectro-polarimeter VEGA installed on the CHARA optical array can be an efficient means of probing the structure and the kinematics of the hot circumstellar gas at sub-AU.} {For the first time, we observed the Herbig~Ae star AB Aur in  the  H$\\alpha$  emission  line, using the  VEGA  low  spectral resolution (R=1700) on two baselines of the array.} {% We computed and calibrated the spectral visibilities of AB~Aur between 610~nm and 700~nm in spectral bands of 20.4~nm. To simultaneously reproduce the line profile and the inferred  visibility around H$\\alpha$, we used a 1-D radiative transfer code (RAMIDUS/PROFILER) that calculates level populations for hydrogen atoms in a spherical geometry and that produces synthetic spectro-interferometric observables.} {We clearly resolved AB Aur in the H$\\alpha$ line and in a part of the continuum, even at the smallest baseline of 34~m. The small P-Cygni absorption feature is indicative of an outflow but could not be explained by a spherical stellar wind model. Instead, it favors a magneto-centrifugal X-disk or disk-wind geometry. The fit of the spectral visibilities from 610 to 700~nm could not be accounted for by a wind alone, so another component inducing a visibility modulation around H$\\alpha$  needed  to  be considered. We thus considered a brightness asymmetry possibly caused by large-scale nebulosity or by the known spiral structures. } {Thanks to the unique capabilities of VEGA, we managed to simultaneously record for the first time a spectrum at a resolution of 1700 and spectral visibilities in the visible range on a target as faint as $m_{V}$ = 7.1. It was possible to rule out a spherical geometry for the wind of AB~Aur and provide realistic solutions to account for the H$\\alpha$ emission compatible with magneto-centrifugal acceleration. It was difficult, however, to determine the exact morphology of the wind because of the surrounding asymmetric ne\\-bulosity. The study illustrates the advantages of optical interferometry and motivates observations of other bright young stars in the same way to shed light on the accretion/ejection processes.} ",
        "introduction": "The class of Herbig Ae/Be (HAEBE) stars gathers the pre-main-sequence stars of intermediate  mass (1.5 $\\leq  M/M_{\\sun} \\leq$ 10),  with spectral types of B to F8, strong emission lines, and the presence of infrared to submillimeter excess flux. Like  their solar-type analogs, the T Tauri objects, HAeBe  objects are known to be surrounded  by protoplanetary disks of gas and dust, responsible for this excess emission. In the case of lower mass stars, the formation is commonly explained by the gravitational collapse of a dusty disk cloud and magnetically controlled accretion via a disk. The formation of more massive stars is still very uncertain. However, determining if similar processes are at work in these objects is extremely important, as it would help to esta\\-blish whether the accretion/ejection mechanism is universal irrespective of the mass of the central body. % The hot gas that can be involved in accretion and  ejection flows  close to  the source and can be  used to  probe the corresponding  physical  conditions  through  its emission  in  lines, especially those in the optical range that have a high diagnostic potential. These phenomena, however, occur in a small region of a few astronomical units (AU) around the star, corresponding,   for   close  star   formation   regions,   to  a   few milliseconds of arc. Such  a small scale can only be resolved with interferometry.  Guided by these considerations, we have attempted for the first time to observe a bright visible Herbig Ae/Be star, AB~Aurigae, with a spectro-interferometer operating at optical wavelengths, the VEGA spectrograph (\\cite{vega}) installed at the CHARA Array (\\cite{chara}). This unique combination of spectral and spatially resolved information allowed us to study the physical conditions for the gas producing the H$\\alpha$ line on sub-AU scales. Interpreting the results, however, required appropriate modeling of the emitted radiation. Using a radiative transfer code, we obtained observational insights on the generation mechanism of winds around AB~Aurigae, giving experimental support to the widely accepted, but not yet tested, theory of magneto-centrifugal acce\\-leration. In Sect.~2, we recall the present knowledge about the target. In Sect.~3, we describe the observations and the data processing. In Sect.~4, we present the radiative transfer code and discuss the wind models adopted to interpret the observed H$\\alpha$ visibilities and spectrum. ",
        "conclusions": "\\begin{figure*}[t] \\centering \\includegraphics[width=9.1cm, angle=0]{14720fg3.pdf} \\includegraphics[width=9.1cm, angle=0]{14720fg4.pdf} \\caption{Best model predictions compared to the VEGA spectrum (left) and spectral  visibilities (right), both  in full circles  with error bars. Dashed lines  correspond to the X-wind model  and full lines to the disk-wind  one. For the squared visibilities,  triangles denote the model prediction including both the wind and the contribution of the asymmetric reflection nebula.} \\label{fig:model-disc-binary} \\end{figure*}  The spectroscopic P-Cygni feature in AB~Aur shows that the line is formed at least in part in a wind. Therefore, we focused on a wind model to constrain the properties of the emitting region. To simultaneously reproduce both the line profile and the spectral visibilities, we have adapted two numerical codes in spherical symmetry for the production and transport of hydrogen (H) lines (\\cite{Rajabi}). The first one, RAMIDUS, calculates the level populations of H in a moving fluid for assigned radial dependencies of the basic wind parameters (density, velocity, tempe\\-rature). The code uses the Sobolev approximation and the escape probability method (\\cite{Castor70}). In cascade, the PROFILER  code calculates synthetic images, line profiles and visibilities of H lines for various geometrical configurations. By selecting limited portions of the spherical winds, we approximated 2-D axisymmetric models. We did not include accretion in the models.  A full spherical wind providing matter along the line of sight can be readily  excluded,  because it  would  produce  too much  blue-shifted absorption  with respect to the observed  P-Cygni  feature that  is shallow but not very deep (normalized flux$\\geq$0.88). Even if the issue of the inclination angle of the AB Aur system is still alive (see Sect. 2), we chose a wind geometry in agreement with a small inclination (i$\\leq40^{\\circ}$), and thus excluded a double cone geometry, reminiscent  of a bipolar  jet, which would lead, for small inclinations, to a line profile with two distinct emission features instead of the observed single peak. Inspired by the magneto-centrifugal (MC) scenario for  the acceleration of  stellar jets, we assume that, on the observed scales (much less than 2-3 AU from the star/disk), the wind occupies a region delimited by a  flattened torus defined by the pale gray region in Fig.~\\ref{fig:model}.  This configuration can simulate       the      base       of       either      a       X-wind (\\cite{Camenzind90}, \\cite{Shu95}) or a disk-wind (hereafter D-wind; \\cite{Ferreira97},   \\cite{KoniglPudritz})   accelerated   along   the magnetic field  lines anchored  in the inner  few AU of  the accretion disk.  In this region the wind has a wide angle aperture, as it has not yet reached the  Alfv\\`en  surface, after  which  magnetic collimation occurs and the wind collimates into a jet. In  Fig.~\\ref{fig:model}, R$_i$=0.02~AU  is the  corotation radius, R$_d$=0.24~AU is the inner radius of the optically thick dusty disk, R$_f$=2.3~AU is  the extension  of the  wind, and $\\alpha $ is the angle between the disk mid-plane and the surface of the wind, constrained to be  $\\leq 60^{\\circ}$ for MC acceleration to occur (\\cite{Spruit96}).  In our X-wind model, all  the material in the flow comes  from the region located very close to the corotation radius and has  a constant  mass outflow rate  $\\dot{\\rm{M}}_w$ through  the torus section, while for the D-wind $\\dot{\\rm{M}}_w$ progressively increases from R$_i$ to about 1~AU, to simulate that the wind originates in an extended region of the disk. The wind velocity is assumed to be $V(r)=V_{0}+V_{f}(1-(R_{\\star}/r)^2)$, where $V_{0} =$ 20 km.s$^{-1}$ is the initial velocity, $r$ the radial distance from the star, $R_{\\star}$ the stellar radius, and $V_{f}$ the terminal velocity.  The density profile is dictated by $V(r)$ and $\\dot{\\rm{M}}_w (r)$ through the mass continuity equation, while the gas temperature profile T$(r)$, which reaches a maximum of 12~000 K, is chosen in accordance with the theory of wind heating illustrated by Shang et al. (1998) and Garcia et al. (2001) for the X-wind and D-wind, respectively. Modeling of the Stark absorption in the extended line wings has been included using a template spectrum from Kurucz~(1979). %  Regarding the modeling of the line profile, both an X-wind and a D-wind could provide a good fit (Fig.~\\ref{fig:model-disc-binary}-left) for the model parameters collected in Table~\\ref{tab:model}. To reproduce at the same time the line width and the absence of a strong blue-shifted absorption, the best predictions for the X-wind have nearly a null inclination ($i$) and a large opening angle ($\\alpha$) of 60$^{\\circ}$. For the D-wind model, since the disk hides part of the high velocity material in the red-shifted wing, the best predictions are obtained with a small inclination (20$^{\\circ}$) and a smaller opening  angle of 35$^{\\circ}$ to avoid  too  much  blue-shifted absorption. Both models lead to a small inclination in agreement with the previous interferometric determinations. As pointed out by  Pogodin (1992), a change in the opening angle of the wind ($\\alpha$) affects the depth of the blue-shifted absorption. Concerning the mass loss rate $\\dot{\\rm{M}}_w$, our best prediction is of the same order as previous determinations from line profile analysis (\\cite{CatalaKunasz}) for the X-wind but is significantly higher for the D-wind. In fact, our various simulations clearly show that the estimate of $\\dot{\\rm{M}}_{\\rm  w}$ is very dependent on the velocity law and the temperature profile considered in the wind model, and on the scale on which the wind is accelerated. Since the geometry of our D-wind model strongly deviates from a spherical wind, the value of $\\dot{\\rm{M}}_{\\rm  w}$ obtained for this model cannot be directly compared to the previous determinations of $\\dot{\\rm{M}}_w$.  \\begin{table}[t] \\centering \\caption{Parameters of  the models, with N$_{\\rm{H},ej}$ the average hydrogen density at the wind footpoints. } % \\label{tab:model} \\begin{tabular}{c c c c c c c} \\hline Model  &   $i$   &  $\\alpha$   &  V$_{\\rm{f}}$ [km/s] & $\\dot{\\rm{M}}_w$ [M$_{\\odot}$/yr]& N$_{\\rm{H},ej}$ [cm$^{-3}$]\\\\ \\hline X-wind & 0$^\\circ$ & 60$^\\circ$  & 410 &  3.7 10$^{-8}$ & 4.3~$10^{10}$ \\\\ D-wind & 20$^\\circ$ & 35$^\\circ$ & 450 &  2.2  10$^{-7}$ & 7.7~$10^{8}$\\\\ \\hline \\end{tabular} \\end{table} \\begin{figure}[t] \\centering \\includegraphics[width=0.35\\textwidth]{14720fg5.pdf} \\caption{Geometry adopted for the empirical X-wind and disk-wind mo\\-dels described in Table~\\ref{tab:model}.  R$_i$, R$_d$, and R$_f$ are the corotation radius, dust sublimation radius, and extension of the wind, respectively.} \\label{fig:model} \\end{figure}  Regarding the modeling of V$^2$, the X-wind  provides only a negligible decrease in H$\\alpha$ ($V^2$ = 0.98) with           respect            to           the           continuum (Fig.~\\ref{fig:model-disc-binary}-right). In this  case, the emission is concentrated  at R$_i$, while it is  much more extended in the D-wind model that produces a larger decrease in $V^2$ (down to 0.89). In both cases, however, the observed visibility variation could not be modeled  with the  wind alone so we considered an additional component producing an extended continuum emission. This component could come from the extended reflection nebulosity around the system, already reported by several authors (\\cite{Gradyb, Fukagawa, MG06}). In practice, we used this secondary source to reproduce the modulation of the visibility in the continuum. Our wind models combined with this additional emission provided a good fit of the visibilities. For instance, Fig.~\\ref{fig:model-disc-binary}-right corresponds to a secondary source separated from the star by 38.2 mas (i.e. 5.5 AU) contributing to 7\\% of the total flux. Because of the large number of free parameters and the fact that we only have measurements along one baseline direction, we found many different configurations for the secondary emission that could fit, for both  wind  models, the  observed visibilities (e.g. a much larger separation of 591.2 mas).  In conclusion, thanks to the unique capabilities of VEGA/CHARA, we managed for the first time to record simultaneously a spectrum at a resolution of 1700 and a visibi\\-lity spectrum in the visible range on a target as faint as $m_{V}$ = 7.1. Despite the limited signal-to-noise ratio of our data, they allowed us to rule out a spherical wind and to make realistic predictions on the nature of the H$\\alpha$ emission region at sub-AU, and our data are compatible with a magneto-centrifugal mechanism for the production of the wind. It was difficult, however, to determine the exact morphology of the wind and disentangle the X-wind and the D-wind because of the extended nebu\\-losity. Moreover, all models neglect the contribution of mass accretion to the H-line emission, so self-consistent models of winds including this process are needed. Complementary observations with larger CHARA baselines are required to place more stringent constraints on the complex environment and the photosphere contribution of AB~Aur that can be resolved in the visi\\-ble with VEGA. We also expect to take advantage of the new instrumental facilities of CHARA to study other bright HAeBe and provide promising insights on the accretion/ejection mechanisms in intermediate-mass young stars."
    },
    "1005/1005.2501_arXiv.txt": {
        "abstract": "We follow the approach of induced--matter theory for a five--dimensional ($5D$) vacuum Brans--Dicke theory and introduce induced--matter and induced potential in four dimensional ($4D$) hypersurfaces, and then employ a generalized FRW type solution. We confine ourselves to the scalar field and scale factors be functions of the cosmic time. This makes the induced potential, by its definition, vanishes, but the model is capable to expose variety of states for the universe. In general situations, in which the scale factor of the fifth dimension and scalar field are~not constants, the $5D$ equations, for any kind of geometry, admit a power--law relation between the scalar field and scale factor of the fifth dimension. Hence, the procedure exhibits that $5D$ vacuum FRW--like equations are equivalent, in general, to the corresponding $4D$ vacuum ones with the same spatial scale factor but a new scalar field and a new coupling constant, $\\tilde{\\omega}$. We show that the $5D$ vacuum FRW--like equations, or its equivalent $4D$ vacuum ones, admit accelerated solutions. For a constant scalar field, the equations reduce to the usual FRW equations with a typical radiation dominated universe. For this situation, we obtain dynamics of scale factors of the ordinary and extra dimensions for any kind of geometry without any \\emph{priori} assumption among them. For non--constant scalar fields and spatially flat geometries, solutions are found to be in the form of power--law and exponential ones. We also employ the weak energy condition for the induced--matter, that gives two constraints with negative or positive pressures. All types of solutions fulfill the weak energy condition in different ranges. The power--law solutions with either negative or positive pressures admit both decelerating and accelerating ones. Some solutions accept a shrinking extra dimension. By considering non--ghost scalar fields and appealing the recent observational measurements, the solutions are more restricted. We illustrate that the accelerating power--law solutions, which satisfy the weak energy condition and have non--ghost scalar fields, are compatible with the recent observations in ranges $-4/3<\\omega\\leq-1.3151$ for the coupling constant and $1.5208\\leq n<1.9583$ for dependence of the fifth dimension scale factor with the usual scale factor. These ranges also fulfill the condition $\\tilde{\\omega}>-3/2$ which prevents ghost scalar fields in the equivalent $4D$ vacuum Brans--Dicke equations. The results are presented in a few tables and figures. ",
        "introduction": "\\indent  Attempts to geometrical unification of gravity with other interactions, using higher dimensions other than our conventional $4D$ space--time, began shortly after invention of the special relativity (\\textbf{SR}). Nordstr{\\o}m was the first who built a unified theory on the base of extra dimensions~\\cite{1}. Tight connection between SR and electrodynamics, namely the Lorentz transformation, led Kaluza~\\cite{2} and Klein~\\cite{3} to establish $5D$ versions of general relativity (\\textbf{GR}) in which electrodynamics rises from the extra fifth dimension. Since then, considerable amount of works have been focused on this idea either using different mechanism for compactification of extra dimension or generalizing it to non--compact scenarios (see e.g. Ref.~\\cite{OverduinWesson1997}) such as Brane--World theories~\\cite{4}, space--time--matter or induced--matter (\\textbf{IM}) theories~\\cite{5} and references therein. The latter theories are based on the Campbell--Magaard theorem which asserts that any analytical $N$--dimensional Riemannian manifold can locally be embedded in an $(N+1)$--dimensional Ricci--flat Riemannian manifold~\\cite{6}. This theorem is of great importance for establishing $4D$ field equations with matter sources locally to be embedded in $5D$ field equations without \\emph{priori} introducing matter sources. Indeed, the matter sources of $4D$ space--times can be viewed as a manifestation of extra dimensions. This is actually the core of IM theory which employs GR as the underlying theory.  On the other hand, Jordan~\\cite{9} attempted to embed a curved $4D$ space--time in a flat $5D$ space--time and introduced a new kind of gravitational theory, known as the scalar--tensor theory. Following his idea, Brans and Dicke~\\cite{10} invented an attractive version of the scalar--tensor theory, an alternative to GR, in which the weak equivalence principle is saved and a non--minimally scalar field couples to curvature. The advantage of this theory is that it is more Machian than GR, though mismatching with the solar system observations is claimed as its weakness~\\cite{Bertotti2003fujiBook}. However, the solar system constraint is a generic difficulty in the context of the scalar--tensor theories~\\cite{banerjee2001Sen2001}, and it does~not necessarily denote that the evolution of the universe, at all scales, should be close to GR, in which there are some debates on its tests on cosmic scales~\\cite{bean2009}.  Although it is sometimes desirable to have a higher dimensional energy--momentum tensor or a scalar field, for example in compactification of extra curved dimensions~\\cite{11}, but the most preference of higher dimensional theories is to obtain macroscopic $4D$ matter from pure geometry. In this approach, some features of a $5D$ vacuum Brans--Dicke (\\textbf{BD}) theory based on the idea of IM theory have recently been demonstrated~\\cite{12}, in where the role of GR as fundamental underlying theory has been replaced by the BD theory of gravitation. Actually, it has been shown that $5D$ vacuum BD equations, when reduced to four dimensions, lead to a modified version of the $4D$ Brans--Dicke theory which includes an induced potential. Whereas in the literature, in order to obtain accelerating universes, inclusion of such potentials has been considered in \\emph{priori} by hand. A few applications and a $D$--dimensional version of this approach have been performed~\\cite{Ponce1and2,12.5}. Though, in Refs.~\\cite{Ponce1and2}, it has also been claimed that their procedure provides explicit definitions for the effective matter and induced potential. Besides, some misleading statements and equations have been asserted in Ref.~\\cite{12}, and hence we have re--derived the procedure in Section $2$. Actually, the reduction procedure of a $5D$ analogue of the BD theory, with matter content, on every hypersurface orthogonal to an extra cyclic dimension (recovering a modified BD theory described by a 4--metric coupled to two scalar fields) has previously been performed in the literature~\\cite{qiang20052009}. However, the key point of IM theories are based on not introducing matter sources in $5D$ space--times.  In addition, recent measurements of anisotropies in the microwave background suggest that our ordinary $4D$ universe should be spatially flat~\\cite{13}, and the observations of Type Ia--supernovas indicate that the universe is in an accelerating expansion phase~\\cite{14}. Hence, the universe should mainly be filled with a dark energy or a quintessence which makes it to expand with acceleration~\\cite{15}. Then after an intensive amount of work has been performed in the literature to explain the acceleration of the universe.  In this work, we explore the Friedmann--Robertson--Walker (\\textbf{FRW}) type cosmology of a $5D$ vacuum BD theory and obtain solutions and related conditions. This model has extra terms, such as a scalar field and scale factor of fifth dimension, which make it capable to present accelerated universes beside decelerated ones. In the next section, we give a brief review of the induced modified BD theory from a $5D$ vacuum space--time to rederive the induced energy--momentum tensor, as has been introduced in Ref.~\\cite{12}, for our purpose to employ the energy density and pressure. In Section~$3$, we consider a generalized FRW metric in the $5D$ space--time and specify FRW cosmological equations and employ the weak energy condition (\\textbf{WEC}) to obtain the energy density and pressure conditions. Then, we probe two special cases of a constant scale factor of the fifth dimension and a constant scalar field. In Section~$4$, we proceed to exhibit that $5D$ vacuum BD equations, employing the generalized FRW metric, are equivalent, in general, to the corresponding vacuum $4D$ ones. This equivalency can be viewed as the main point within this work which distinguishes it from Refs.~\\cite{12,Ponce1and2}. In Section~$5$, we find exact solutions for flat geometries and proceed to get solutions fulfilling the WEC while being compatible with the recent observational measurements. We also provide a few tables and figures for a better view of acceptable range of parameters. Finally, conclusions are presented in the last section. ",
        "conclusions": "\\indent  Analogous to the approach of IM theories, one can consider the BD gravity as the underlying theory. Hence, extra geometrical terms, coming from the fifth dimension, are regarded as an induced--matter and induced potential. We have followed, with some corrections, the procedure of Ref.~\\cite{12} for introducing the induced potential and have employed a generalized FRW type solution for a $5D$ vacuum BD theory. Hence, the scalar field and scale factors of the $5D$ metric can, in general, be functions of the cosmic time and the extra dimension. However, for simplicity, we have assumed the scalar field and scale factors to be only functions of the cosmic time, where this makes the induced potential, by its definition, vanishes.  We then have revealed that in general situations, in which the scale factor of the fifth dimension and scalar field are~not constants, the $5D$ equations, for any kind of geometry, admit a power--law relation between the scalar field and scale factor of the fifth dimension. Hence, the procedure exhibits that $5D$ vacuum FRW--like equations are equivalent, in general, to the corresponding $4D$ vacuum ones with the same spatial scale factor but a new (or modified) scalar field and a new coupling constant. This equivalency can be viewed as the distinguished point of this work from Refs.~\\cite{12,Ponce1and2}. Indeed, through investigating the $5D$ vacuum FRW--like equations, we have shown that its equivalent $4D$ vacuum equations admit accelerated scale factors, contrary to what one may have expected from a vacuum space--time. Conclusions of the complete investigation of the induced $4D$ equations are as follows.  Following our investigations for cosmological implications, we have shown that for the special case of a constant scale factor of the fifth dimension, the $5D$ vacuum FRW--like equations reduce to the corresponding equations of the usual $4D$ vacuum BD theory, as expected. In the special case of a constant scalar field, the action reduces to a $5D$ Einstein gravitational theory and the equations reduce to the usual FRW equations with a typical radiation dominated universe. For this situation, we also have obtained dynamics of scale factors of the ordinary and extra dimensions for any kind of geometry without any \\emph{priori} assumption among them. Solutions predict a limited life time for closed geometries and unlimited one for flat and open geometries. A typical time evolutions of scale factors correspond to closed, flat and open geometries have been illustrated in Fig.~$0$.  Then, we have focused on spatially flat geometries and have obtained exact solutions of scale factors and scalar field. Solutions are found to be in the form of power--law and exponential ones in the cosmic time. We also have employed the WEC for the induced--matter of the $4D$ modified BD gravity, that gives two conditions (\\ref{49.1}) and (\\ref{49.2}). We then have pursued properties of these solutions and have indicated mathematically and physically acceptable ranges of them, and the results have been presented in a few tables and figures.  All types of solutions fulfill the WECs in different ranges, where the exponential solutions are more restricted. The solutions fulfilling the WEC (\\ref{49.1}) have negative pressures, but the figures illustrate that for the power--law results there are decelerating solutions beside accelerating ones. For this condition, both $\\rho_{_{\\rm BD}}$ and $|p_{_{\\rm BD}}|$ decrease with the cosmic time, but the extra dimension grows. On the other hand, the solutions satisfying the WEC (\\ref{49.2}) have positive pressures, where the power--law results accept accelerating solutions in addition to decelerating ones. For this condition, again decreasing energy density and pressure with the time can occur for some solutions, however all with shrinking extra dimension. The homogeneity between the extra dimension and the usual spatial dimensions, i.e. $b\\propto a$, can take place in the solutions, but for the power--law ones the WECs exclude it.  By considering non--ghost scalar fields and appealing the recent observational measurements, the solutions have been more restricted. Actually, we have illustrated that the accelerating power--law solutions, which satisfy the WEC and have non--ghost scalar fields, are compatible with the recent observations in ranges $-4/3<\\omega\\leq-1.3151$ for the BD coupling constant and $1.5208\\leq n<1.9583$ for dependence of the fifth dimension scale factor with the usual scale factor. These ranges also fulfill the condition $\\tilde{\\omega}>-3/2$ which prevents ghost scalar fields in the equivalent $4D$ vacuum BD equations. Incidentally, this range is more restricted than the one obtained in Ref.~\\cite{Ponce1and2}, i.e. $-1.5<\\omega<-1$, where the difference may have been caused by the distinct definition of the induced potential in two approaches of Ref.~\\cite{12} and Ref.~\\cite{Ponce1and2}. However, we should remind that it has also been shown~\\cite{Freund1982} that the WEC, for $5D$ space--times, requires $-4/3\\leq\\omega$, in which no other experimental evidences have been considered."
    },
    "1005/1005.1502_arXiv.txt": {
        "abstract": "{Herschel PACS and SPIRE images have been obtained of NGC 6720 (the Ring Nebula). This is an evolved planetary nebula with a central star that is currently on the cooling track, due to which the outer parts of the nebula are recombining. From the PACS and SPIRE images we conclude that there is a striking resemblance between the dust distribution and the H$_2$ emission, which appears to be observational evidence that H$_2$ forms on grain surfaces. We have developed a photoionization model of the nebula with the Cloudy code which we used to determine the physical conditions of the dust and investigate possible formation scenarios for the H$_2$. We conclude that the most plausible scenario is that the H$_2$ resides in high density knots which were formed after the recombination of the gas started when the central star entered the cooling track. Hydrodynamical instabilities due to the unusually low temperature of the recombining gas are proposed as a mechanism for forming the knots. H$_2$ formation in the knots is expected to be substantial after the central star underwent a strong drop in luminosity about one to two thousand years ago, and may still be ongoing at this moment, depending on the density of the knots and the properties of the grains in the knots.} ",
        "introduction": "Grains play an important role in many environments, including planetary nebulae (PNe), because of extinction, photoelectric heating, their influence on the charge and ionization balance of the gas, as catalysts for grain-surface chemical reactions (e.g.\\ H$_2$ formation), and as seeds for freeze-out of molecules. Previous satellite missions such as IRAS, ISO, Spitzer, and AKARI have allowed us to study the dust in PNe, but unfortunately the angular resolution of these instruments was too low to get detailed information on the spatial distribution of the dust. This has now changed with the launch of Herschel, which allows us to study the spatial structures in unprecedented detail. In this paper we will do this for NGC 6720 (M57, the Ring nebula) to study H$_2$ formation. NGC 6720 is an evolved, oxygen-rich bipolar nebula seen nearly pole-on. The nebula is optically thick to ionizing radiation in most directions, but optically thin in the polar regions \\citep[hereafter OD07]{OD07}. Molecules such as H$_2$ and CO have been detected \\citep{Be78, Hu86}. The central star has exhausted hydrogen shell burning and is now on the cooling track. As a result the outer halo is recombining and re-ionization of the innermost recombined material due to expansion of the nebula has just started (OD07). This object is very similar to the Helix nebula, which seems to be further advanced along the same evolutionary path (OD07). In Sect.~\\ref{obs} we will describe the observations, and in Sect.~\\ref{origin} we will discuss various scenarios for the formation of H$_2$. ",
        "conclusions": ""
    },
    "1005/1005.4782_arXiv.txt": {
        "abstract": "In this paper we present results of an in depth analysis of the High Mass X-ray Binary \\gx~ using \\ig~ data. The source has a double-peaked orbital lightcurve, with flares close to orbital phases 0.5 and 0.9. In addition to archival data three dedicated observations with a total exposure of 200~ksec were performed to observe the pre-periastron flare of the source when the flux from the source is maximal. We obtained broad-band (2-100 keV) pulse phase resolved spectra. We discuss the spectral models suitable for spectrum description and the impact of the model choice on CRSF energy pulse phase dependence. The history of the pulse period since the launch of \\ig~ is determined and the spin down trend, established by RXTE \\citep{Pravdo:2001p360,Kreykenbohm:2004p155} and BeppoSax \\citep{LaBarbera:2005p156} observations is confirmed. We update orbital ephemeris of the system published previously by \\cite{Koh:1997p138}, based on time arrival analysis of archival data. The updated value of orbital period 41.482(2) is found to be consistent with the period obtained from the RXTE ASM lightcurve of the source. ",
        "introduction": "% \\label{sec:introduction} \\gx~ (also known as 4U 1223-62) is a High Mass X-ray Binary system, consisting of a Neutron star orbiting an early B-type optical companion \\Wr~. The neutron star is a $\\sim$680s X-ray pulsar \\citep{White:1976p993}, accreting from a dense wind of the optical companion. The wind mass loss rate of the optical component is one of the highest known in the Galaxy ($\\dot{M}\\sim 10^{-5} M_{\\odot}\\cdot yr^{-1}$). Taking into account very low ($v_{0}\\sim 300~km\\cdot s^{-1}$, \\cite{Kaper:2006p1357}) terminal speed of the wind , the accretion rate shall be high enough to explain observed luminosity of order $L_x\\sim10^{37} ~erg\\cdot s^{-1}$. Distance to the source is estimated between 1.8~kpc \\citep{Parkes:1980p1360} and 5.3~kpc \\citep{Kaper:1995p151} depending on spectral classification of \\Wr, with latest estimates close to 3~kpc \\citep{Kaper:2006p1357}. We will adopt the later value through out the paper. The source exhibits regular X-ray flares just 1-2d before periastron passage (orbital phase $\\sim0.9$). There's also a sign of other flare at orbital phase $\\sim0.5$ \\citep{Koh:1997p138}. A number of hypotheses were proposed to explain observed orbital lightcurve, including circumstellar disk \\citep{Koh:1997p138} as well as presence of quasi-stable accretion stream \\citep{Leahy:2008p358}. In this paper we focus mainly on studying properties of the source during pre-periastron flare, there brightness of the source is highest, so spectral properties can be studied with high statistics.  The spectrum of the source is quite complicated and feature rich. Lower energy part of spectra requires use of partial covering models for absorption with optional reflection component. An iron line of complicated shape is observed at $\\sim6.5$ keV \\citep{Watanabe:2003p357}. High energy cutoff at $\\sim20keV$ together with deep and broad cyclotron resonance scattering feature at $\\sim30-45keV$ are main features at higher energies. ",
        "conclusions": "\\label{sec:conclusions}  \\pagebreak \\footnotesize"
    },
    "1005/1005.4057_arXiv.txt": {
        "abstract": "Observations of the frequency dependence of the global brightness temperature of the redshifted 21 cm line of neutral hydrogen may be possible with single dipole experiments.  In this paper, we develop a Fisher matrix formalism for calculating the sensitivity of such instruments to the 21 cm signal from reionization and the dark ages.  We show that rapid reionization histories with duration $\\Delta z\\lesssim 2$ can be constrained, provided that local foregrounds can be well modelled by low order polynomials.  It is then shown that observations in the range $\\nu=50-100{\\,\\rm MHz}$ can feasibly constrain the \\lya and X-ray emissivity of the first stars forming at $z\\sim15-25$, provided that systematic temperature residuals can be controlled to less than 1 mK.  Finally, we demonstrate the difficulty of detecting the 21 cm signal from the dark ages before star formation. ",
        "introduction": "\\label{sec:intro} The transition of the Universe from the dark ages following hydrogen recombination through to the epoch of reionization remains one of the least constrained frontiers of modern cosmology.  Observing the sources responsible for heating and ionizing the intergalactic medium (IGM) at redshifts $z\\gtrsim6$ pushes current observational techniques to the limit.  Plans are underway to construct low-frequency radio telescopes, such as LOFAR\\footnote{http://www.lofar.org/}, MWA\\footnote{http://www.MWAtelescope.org/}, PAPER\\footnote{\\citet{parsons2009}}, and SKA\\footnote{http://www.skatelescope.org/}, to observe the red-shifted 21 cm line of neutral hydrogen.  These experiments aim to map the state of the intergalactic medium via tomographic observations of 3D fluctuations in the 21 cm brightness temperature.  A simpler and significantly lower cost alternative to this would be measurements of the global 21 cm signal integrated over the sky \\cite{mmr1997,shaver1999,sethi2005}, which can be achieved by single dipole experiments like EDGES \\cite{bowman2007edges} or CoRE \\cite{chippendale2005}.  Although such experiments are today in their infancy, their potential is large.    In this paper, we explore the potential for these global sky experiments to measure the 21 cm signal and constrain the high redshift Universe.  We may draw a historical analogy with the Cosmic Background Explorer (COBE), whose FIRAS instrument measured the blackbody spectrum of the cosmic microwave background (CMB) \\cite{mather1994} while the DMR instrument measured the level of temperature fluctuations \\cite{smoot1992}.  The precise measurement of a $T_{\\rm CMB}=2.726{\\,\\rm K}$ blackbody spectrum placed tight constraints on early energy injection, since no Compton-$y$ or $\\mu$-distortion were seen, and provided important evidence confirming the big bang paradigm.  The detection of angular fluctuations paved the way for more sensitive experiments such as BOOMERANG \\cite{lange2001} and WMAP \\cite{spergel2003}, which provided precision measurements of the CMB acoustic peaks.  While, at the moment, attention is focussed on experiments designed to measure 21 cm fluctuations, it is important not to neglect the possibility of measuring the global signal.  The evolution of the 21 cm signal is driven primarily by the amount of neutral hydrogen and the coupling between the 21 cm spin temperature and the gas temperature.  It is able to act as a sensitive thermometer when the IGM gas temperature is less than the CMB temperature placing constraints on energy injection that leads to heating.  For example, the first black holes to form generate X-rays, which heat the gas.  More exotic processes such as annihilating dark matter might have also been important.  Additionally, energy injection in the form of \\lya production modifies the strength of the coupling.  This provides a way of tracking star formation, which will be the dominant source of \\lya photons.  As we show, the spectral structure of the 21 cm signal is much richer than that of a blackbody so that many things can be learnt about the early Universe. Given the uncertainties, we develop a model approach based upon those physical features most likely to be present.  The single most important factor determining the sensitivity of dipoles to astrophysics will be their ability to remove galactic foregrounds \\cite[e.g.][]{oh2003,dimatteo2004}.  Exploitation of spectral smoothness to remove foregrounds by fitting low order polynomials is key to avoiding throwing the signal away with the foreground.  To quantify this, we develop a simple Fisher matrix formalism and validate it against more detailed numerical parameter fitting.  This provides us with a way of quantitatively addressing the ability of global 21 cm experiments to constrain reionization and the astrophysics of the first galaxies \\cite{loeb2010book}.  Similar work on the subject \\cite{sethi2005} ignored the influence of foregrounds limiting its utility considerably.  Much of the power of this technique stems from the limitations of other observational probes.  While next generation telescopes such as JWST\\footnote{http://www.jwst.nasa.gov/}, GMT\\footnote{http://www.gmto.org/}, EELT\\footnote{http://www.eso.org/sci/facilities/eelt/} or TMT\\footnote{http://www.tmt.org/} may provide a glimpse of the Universe at $z\\gtrsim12$ they peer through a narrow field of view and are unlikely to touch upon redshifts $z\\gtrsim20$.  As we will show, 21 cm global experiments could potentially provide crude constraints on even higher redshifts at a much lower cost.  The structure of this paper is as follows. In \\S\\ref{sec:physics}, we begin by describing the basic physics that drives the evolution of the 21 cm global signature and drawing attention to the key observable features.  We follow this in \\S\\ref{sec:foreground} with a discussion of the foregrounds, which leads into our presenting a Fisher matrix formalism for predicting observational constraints in \\S\\ref{sec:fisher}.  In \\S\\ref{sec:reion} and \\S\\ref{sec:astro} we apply this formalism to the signal from reionization and the first stars, respectively.  After a brief discussion in \\S\\ref{sec:darkages} of the prospects for detecting the signal from the dark ages before star formation, we conclude in \\S\\ref{sec:conclude}.  Throughout this paper where cosmological parameters are required we use the standard set of values $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $\\Omega_b=0.046$, $H=100h\\,\\rm{km\\,s^{-1}\\,Mpc^{-1}}$ (with $h=0.7$), $n_S=0.95$, and $\\sigma_8=0.8$, consistent with the latest measurements \\citep{komatsu2009}. ",
        "conclusions": "\\label{sec:conclude}  Observations of the redshifted 21 cm line potentially provide a new window into the high redshift Universe.  Detecting this signal in the presence of large foregrounds is challenging and it is important to explore all avenues for exploiting the signal.  In this paper, we have focussed upon the possibility of using single dipole experiments to observe the all-sky 21 cm signal, in contrast to the 21 cm fluctuations targeted by MWA, LOFAR, PAPER, and SKA.  Experiments targeting this global signal are in their infancy.  We emphasise that instruments built from a few dipoles targeting the global 21 cm signal can be several orders of magnitude cheaper to build than interferometers targeting the fluctuations.  Their scientific return will be similarly less, but at this stage where we know so little about the first sources, even that little is extremely valuable.  As we have outlined in this paper, the 21 cm signal generated by astrophysical processes has a well defined form, although the input parameters are only poorly understood.  We have demonstrated that, at the level of our current knowledge, describing the Galactic foregrounds requires at least a 3rd order polynomial.  At this level, we are able to remove the foregrounds to the sub-mK level, although in practice this procedure may be more complicated.  In order to characterise the sensitivity of these experiments to the signal, we developed a Fisher matrix formalism and validated it against more numerical fitting of the model parameters.  This Fisher matrix approach allows rapid calculations of the experimental sensitivity and appears to reproduce more detailed calculations very well.  Having developed this formalism we applied it to the signal from reionization and the epoch of the first stars.  Using a toy model of reionization, we demonstrated that EDGES-like experiments should be capable of constraining rapid reionization histories with $\\Delta z\\lesssim2$.  More promisingly, these experiments can rule out a wide variety of astrophysical models for the signal from the first stars where the evolution of the spin temperature is important.  We used a straightforward fitting form for the signal based upon the positions of the turning points and showed that these features could be constrained, with the deepest absorption trough providing the best observational target.  Finally, we briefly explored the possibility of detecting the absorption feature present before star formation began.  The increased foreground brightness at low frequencies make it very difficult to constrain this feature and will require long integration times and more sophisticated methods of foreground removal.  This paper represents a first serious look at the prospects for using global measurements of the 21 cm signal to constrain astrophysics.  As a result, there are a number of places where future work might improve upon our calculations.  These include investigating the effects of finite sky coverage, incorporating an arbitrary instrumental frequency response, and allowing for the removal of frequency channels corrupted by terrestrial radio interference."
    },
    "1005/1005.0629.txt": {
        "abstract": "The Advanced Camera for Surveys (ACS) and the Near Infrared Camera and Multi-Object Spectrometer  (NICMOS) have been used to obtain new {\\em Hubble Space Telescope} images of NGC~4038/4039 (``The Antennae''). These new observations allow us to better differentiate compact star clusters from individual stars, based on both size and color. We use this ability to extend the cluster luminosity function by approximately two magnitudes over our previous WFPC2 results, and find that it continues as a single power law, $dN/dL \\propto L^{\\alpha}$ with $\\alpha = -2.13 \\pm 0.07$, down to the observational limit of $\\Mv\\approx -7$. Similarly, the mass function is a single power law $dN/dM \\propto M^{\\beta}$ with $\\beta = -2.10 \\pm 0.20$ for clusters with ages  $<\\!3\\times10^8$~yr, corresponding to lower mass limits that range from $10^4$ to $10^5\\,\\msun$, depending on the age range of the subsample. Hence the power law indices for the luminosity and mass functions are essentially the same. The luminosity function for intermediate-age clusters (i.e., $\\sim$100--300 Myr old objects found in the loops, tails, and outer areas) shows no bend or turnover down to $\\Mv\\approx -6$, consistent with relaxation-driven cluster disruption models which predict the turnover should not be observed until $\\Mv \\approx -4$. An analysis of individual $\\sim$0.5-kpc sized areas over diverse environments shows good agreement between values of $\\alpha$ and $\\beta$, similar to the results for the total population of clusters in the system. There is tentative evidence that the values of both $\\alpha$ and $\\beta$ are flatter for the youngest clusters in some areas, but it is possible that this is caused by observational biases. Several of the areas studied show evidence for age gradients, with somewhat older clusters appearing to have triggered the formation of younger clusters. The area around Knot~B is a particularly interesting example, with an $\\sim$10--50~Myr old cluster of estimated mass $\\sim\\!10^6\\,\\Msun$ having apparently triggered the formation of several younger, more massive (up to $5\\times10^6\\,\\Msun$) clusters along a dust lane. A comparison with new NICMOS observations reveals that only $16 \\pm 6$\\% of the IR-bright clusters in the Antennae are still heavily obscured, with values of $\\Av > 3$ mag. ",
        "introduction": "The Antennae (NGC 4038/39) are the youngest and nearest example of a pair of merging disk galaxies in the Toomre (1977) sequence. As such, they have been observed in virtually every wavelength regime: radio (e.g., Wilson \\etal\\ 2000, 2003; Neff \\& Ulvestad 2000; Hibbard \\etal\\ 2001; Zhu, Seaquist \\& Kuno 2003; Schulz \\etal\\ 2007), infrared (e.g., Mirabel \\etal\\ 1998; Kassin \\etal\\ 2003; Wang \\etal\\ 2004; Brandl \\etal\\ 2005, 2009; Mengel \\etal\\ 2005; Clark \\etal\\ 2007; Rossa \\etal\\ 2007; Gilbert \\& Graham 2007; Mengel \\etal\\ 2008), optical (Rubin, Ford, \\& D'Odorico 1970; Whitmore \\& Schweizer 1995; Whitmore \\etal\\ 1999; Kassin \\etal\\ 2003; Bastian \\etal\\ 2006), ultraviolet (e.g., Hibbard \\etal\\ 2005), and X-ray (e.g., Fabbiano \\etal\\ 2003, 2004; Metz \\etal\\ 2004; Zezas \\etal\\ 2006, 2007).  One important result has been the discovery that much of the star formation occurs in the form of massive compact star clusters (Whitmore \\& Schweizer 1995; Whitmore \\etal\\ 1999), often referred to as ``super star clusters.'' Some of these clusters have properties expected of young globular clusters (i.e., mass, effective radius, etc.).  Hence, it may be possible to study the formation (and disruption, see Fall, Chandar, \\& Whitmore 2005, hereafter FCW05; Whitmore, Chandar, \\& Fall 2007, hereafter WCF07; and Fall, Chandar, \\& Whitmore 2009, hereafter FCW09) of globular clusters in the local universe rather than trying to determine how they formed some 13 Gyr ago.  One of the primary challenges for previous observations of the Antennae was the difficulty of distinguishing faint young clusters from individual young stars, which can be as luminous as $\\Mv \\approx -9$. The Advanced Camera for Surveys, with its superior spatial resolution over a wider field of view than the Wide Field and Planetary Camera 2 (WFPC2), provides an opportunity to improve this situation. The ability to better distinguish faint clusters from stars provides several new science opportunities.  The first is an extension of the cluster luminosity and mass functions to lower values. A large number of studies have found that the young cluster luminosity function (LF) in a variety of star-forming galaxies is a power law of the form $\\phi(L)dL \\propto L^{\\alpha} dL$ with a value of $\\alpha\\approx -2$ (e.g., see compilations in Whitmore 2003; Larsen 2005, 2006; de Grijs \\& Parmentier 2007). However, recently there have been claims (e.g., Fritze-von Alvensleben 2004; de Grijs \\& Parmentier 2007; Anders \\etal\\ 2007) that the initial cluster luminosity function has a turnover at faint magnitudes, and is closer to the Gaussian form found for old globular clusters. The ability to push the LF about 2--3 mag fainter provides a more stringent test of the power-law nature of the initial cluster LF.  Areas in the outskirts of the galaxy provide an especially sensitive test, since their low background levels allow cluster detections to fainter magnitudes. Most of the clusters in these tails, loops, and outer areas have intermediate ages (approximately 100--300 Myr) and have, therefore, survived the initial stages of cluster disruption or ``infant mortality'' associated with the removal of interstellar material (see Hills 1980; Lada, Margulis \\& Dearborn 1984; Boily \\& Kroupa 2003;  Bastian \\& Goodwin 2006; and Fall \\etal\\ 2010 for theoretical discussions and Lada \\& Lada 2003, FCW05, and WCF07 for observational evidence).  Several studies have also found that the young cluster mass function (MF) in a variety of galaxies is a power law of the form $\\phi(M)dM \\propto M^{\\beta} dM$ with a value of $\\beta\\approx -2$ (e.g., Zhang \\& Fall 1999 in the Antennae; Hunter \\etal\\ 2003 in the Magellanic Clouds; Bik \\etal\\ 2003 in M51; and Chandar \\etal\\ 2010 in M83). While the LF and MF are sometimes considered to be interchangeable, this is not true  for populations of clusters that include a wide range of ages, and hence mass-to-light ratios. In fact, the observed similarity in the values of $\\alpha$ and $\\beta$ can be used as indirect evidence that the age and mass distributions are independent, at least up to ages of 300 Myr (e.g., Fall 2006).  This independence has been established more directly in FCW05, and can be expressed explicitly by the fact that the bivariate mass-age distribution $g(M,\\tau)$ can be approximated by the product of the univariate mass and age distribution:  $g(M,\\tau)\\approx M^{\\beta} \\tau^{\\gamma}$  (FCW05, WCF07, FCW09).  There have also been claims for several galaxies, including the Antennae, that there is a cutoff to the upper mass with which clusters can form (e.g., Gieles \\etal\\ 2006; Bastian 2008; Larsen 2009), and that there are bends in the age distributions of some galaxies (e.g., Boutloukos \\& Lamers 2003, and Lamers \\etal\\ 2005). However, based on our WFPC2 data, we found that both the mass function and the age distribution are simple power laws (see FCW09 for discussion). We suspect that these apparent discrepancies are due to selection effects in other studies, in particular to the exclusion of crowded regions  which contain most of the young clusters.  In this paper we examine whether or not there is evidence for such bends or cutoffs from our new ACS data set.  A second new science opportunity enabled by the high-resolution ACS images is to dissect the contents, both clusters and individual stars, in some of the spectacular star forming knots found in the Antennae. We will be able to assess whether some of the universal trends found for star-forming galaxies, for example the correlation between the absolute visual magnitude of the brightest cluster and the total number of clusters (e.g., Whitmore 2003), holds or not in individual sub-kpc regions of the Antennae. Age and mass estimates for the clusters in these regions will allow us to assess how star formation proceeds in these extreme environments, and whether there is evidence for sequential star formation.  In this paper, we adopt a distance modulus of 31.71 mag for the Antennae, as found by Schweizer \\etal\\ (2008) based on the Type Ia supernova 2007sr. This modulus corresponds to a distance of 22~Mpc for a Hubble constant of $H_0 = 72$ km s$^{-1}$ Mpc$^{-1}$. Note that this distance is slightly further than that adopted in our earlier papers (19.2 Mpc), and quite different from the shorter distance of 13.3 Mpc found by Saviane \\etal\\ (2008), based on an apparent misidentification of the tip of the red giant branch (see Schweizer \\etal\\ 2008 for details). At a distance of 22 Mpc, 1\\arcsec\\ is equivalent to 107 pc, and one $0\\farcs05$ ACS/WFC pixel covers 5.33 pc.  The remainder of this paper is organized as follows: \\S2 describes the observations and photometric reductions; \\S3 gives details of our procedure to estimate the ages and masses of the clusters; and \\S4  summarizes our method for separating stars and clusters (with more details included in Appendix~A). We compare the 50 most luminous, massive, and IR-bright clusters in \\S5; examine the LF of star clusters in the Antennae in \\S6, and then the mass and age distributions in \\S7. In \\S8 we perform a detailed analysis of the cluster populations in star-forming knots and other regions, and \\S9 discusses patterns of star formation, including triggered star formation. A summary is provided in \\S10. Future papers will include a more detailed study of the age, mass, and size distributions of the clusters (Chandar \\etal, in prep.) and will incorporate our NICMOS observations of the Antennae more fully (Rothberg \\etal, in prep). ",
        "conclusions": "The Advanced Camera for Surveys on {\\em HST\\/} was used to obtain high-resolution images of NGC~4038/4039, which allow us to better differentiate between individual stars and compact star clusters.  We have used this improved ability to extend the cluster LF by approximately 2--3 mag at its faint end, when compared to the previous WFPC2 photometry. In addition, the Near Infrared Camera and Multi-Object Spectrometer (NICMOS) was used to help study the clusters still largely embedded in dust. Following are our main results:  1. The cluster luminosity function continues as a single power law, $dN/dL \\propto L^{\\alpha}$, with $\\alpha=-2.13 \\pm 0.07$ down to the current observational limit of $\\Mv \\approx -7$. This value is uncorrected for extinction and based on variable-binning fits. Similarly, the cluster mass function is a single power law $dN/dM \\propto M^{\\beta}$ with $\\beta = -2.10 \\pm 0.20 $ for clusters with ages  $<\\!3\\times10^8$~yr, corresponding to lower mass limits that range from $10^4$ to $10^5\\,\\msun$, depending on the age range of the subsample. Hence the power-law indices for the luminosity and mass functions are essentially the same.  2. The LF for intermediate-age clusters (e.g., $\\sim$200 Myr old clusters found in the loops, tails, and outer areas) does not show any hint of a turnover down to $\\Mv \\approx -6$. This is consistent with relaxation-driven cluster disruption models, which predict the turnover should not be observed until $\\Mv \\approx -4$ for this age range.  3. The brightest individual stars in the Antennae reach absolute visual   magnitudes of $\\Mv \\approx -9.5$, a limit compatible with the so-called Humphreys--Davidson (1979) limit observed in the Milky Way and nearby galaxies (see Appendix A for details).  4. Our photometrically determined cluster ages are in good agreement with spectroscopically determined ages from Bastian \\etal\\ (2009) for 13 of the 15 clusters in common.  In the other two cases, spatial resolution effects are likely to be responsible for the age differences, since the ground-based spectroscopic observations contain light from a much larger area than our \\textit{HST} measurements.  5. An area-by-area analysis shows a strong correlation between the absolute magnitude of an area's brightest cluster, \\Mv(brightest), and the (logarithmic) number of clusters brighter than $\\Mv = -9$ in that area, $\\log N$.  The \\Mv(brightest) vs.\\ $\\log N$ diagram for all areas of the Antennae studied is essentially identical to the corresponding diagram for entire galaxies, supporting the universality of this observed trend on a scale of a few hundred pc and larger. In addition, the reddening vs.\\ age correlation found for individual clusters by Whitmore \\& Zhang (2002), Mengel \\etal\\ (2005), and Bastian \\etal\\ (2005) has been confirmed for the area-by-area analysis as well.  6. The same area-by-area analysis also reveals a tentative discovery of a new relationship between the power-law exponent $\\alpha$ of the LF and the median age of clusters in each area, in the sense that younger clusters have flatter LFs (i.e., values of $\\alpha\\/ \\approx -1.8$ for the youngest star- and cluster-forming knots and $-$2.1 for the outer areas).  However, it is possible that this is caused by observational biases.  The power-law indices for the LFs and MFs are in good agreement, and are similar to the results for the total population of clusters in the galaxy.  7. Using our age and mass estimates of clusters in the Antennae, we have    found evidence of sequential cluster formation in several regions. Often there is a massive, somewhat older cluster which appears to have triggered the formation of younger, lower-mass satellite clusters. The area around Knot~B is a particularly dramatic example of triggered star cluster formation, with evidence for one massive ($\\sim$10$^6\\,\\Msun$) cluster,  with an age in the range 10--50~Myr, having triggered the formation of several younger massive clusters (up to $5\\times 10^6$\\,\\Msun) in the direction of a strong dust lane, with an amplification factor of $\\sim$8$\\times$ in the total mass of triggered clusters.  8. NICMOS observations show that $16 \\pm 6$\\% of the clusters in the Antennae are still largely embedded in their dust cocoons, with values of $3 \\la \\Av \\la 7$ mag. However, the top 50 IR-bright clusters {\\it all\\/} have optical counterparts in the much deeper ACS observations, hence the effect on our optically-based catalog derived from deep ACS images is minimal.  In conclusion, the results presented here confirm and extend those based on previous WFPC2 data (e.g., Whitmore \\etal\\ 1999, Zhang \\& Fall 1999, Whitmore \\& Zhang 2002, FCW05, WCF07, FCW09). In addition, similarities between our results for the Antennae  and those for a growing number of other star-forming galaxies (e.g., Mora \\etal\\ 2009)  support the universality model for the formation and disruption of star clusters."
    },
    "1005/1005.0611.txt": {
        "abstract": "{Free fall has signed the greatest markings in the history of physics through the leaning Pisa tower, the Cambridge apple tree and the Einstein lift. The perspectives offered by the capture of stars by supermassive black holes are to be cherished, because the study of the motion of falling stars will constitute a giant step forward in the understanding of gravitation in the regime of strong field. After an account on the perception of free fall in ancient times and on the behaviour of a gravitating mass in Newtonian physics, this chapter deals with last century debate on the repulsion for a Schwarzschild black hole and mentions the issue of an infalling particle velocity at the horizon. Further, black hole perturbations and numerical methods are presented, paving the way to the introduction of the self-force and other back-action related methods. The impact of the perturbations on the motion of the falling particle is computed via the tail, the back-scattered part of the perturbations, or via a radiative Green function. In the former approach, the self-force acts upon the background geodesic; in the latter, the geodesic is conceived in the total (background plus perturbations) field. Regularisation techniques (mode-sum and Riemann-Hurwitz $z$ function) intervene to cancel divergencies coming from the infinitesimal size of the particle. An account is given on the state of the art, including the last results obtained in this most classical problem, together with a perspective encompassing future space gravitational wave interferometry and head-on particle physics experiments. As free fall is patently non-adiabatic, it requires the most sophisticated techniques for studying the evolution of the motion. In this scenario, the potential of the self-consistent approach, by means of which the background geodesic is continuously corrected by the self-force contribution, is examined. } ",
        "introduction": "The two-body problem in general relativity remains one of the most interesting problems, being still partially unsolved. Specifically, the free fall, one of the eldest and classical problems in physics, has characterised the thinking of the most genial developments and it is taken as reference to measure our progress in the knowledge of gravitation. Free fall contains some of the most fundamental questions on relativistic motion. The mathematical simplification, given by the reduction to a 2-dimensional case, and the non-likelihood of an astrophysical head-on collision should not throw a shadow on the merits of this problem. Instead, it may be seen as an arena where to explore part of the relevant features that occur to general orbits, e.g. the coupling between radial and time coordinates.  Although it is easily argued that radiation reaction has a modest impact on radial fall due to the feebleness of cumulative effects (anyhow, in case of high or even relativistic - a fraction of $c$ - initial velocity of the falling particle, it is reasonable to suppose a non-modest impact on the waveform and possibly the existence of a signature), it would be presumptuous to consider free fall simpler than circular orbits, or even elliptic orbits if in the latter adiabaticity may be evoked. Adiabaticity has been variously defined in the literature, but on the common ground of the secular effects of radiation reaction occurring on a longer time scale than the orbital period. One definition refers to the particle moving anyhow, although radiating, on the background geodesic (local small deviations approximation), of obviously no-interest herein; another, currently debated for bound orbits, to the secular changes in the orbital motion being stemmed solely by the dissipative effects (radiative approximation); the third to the radiation reaction time scale being much longer than the orbital period (secular approximation), which is a rephrasing of the basic assumption.  But in radial fall such an orbital period doesn't exist. And as the particle falls in, the problem becomes more and more complex. In curved spacetime, at any time the emitted radiation may backscatter off the spacetime curvature, and interact back with the particle later on. Therefore, the instantaneous conservation of energy is not applicable and the momentary self-force acting on the particle depends on the particle's entire history. There is an escape route, though, for periodic motion. But energy-momentum balance can't be evoked in radial fall, lacking the opportunity of any adiabatic averaging. The particle reaction to its radiation has thus to be computed and implemented immediately to determine the effects on the subsequent motion. It is a no-compromise analysis, without shortcuts. Thus, the computation and the application of the back-action all along the trajectory and the continuous correction of the background geodesic, it is the only semi-analytic way to determine motion in non-adiabatic cases. And once this self-consistent approach shall be mastered for radial infall, where simplification occurs for the two-dimensional nature of the problem, it shall be applicable to generic orbits.  It is worth reminding that the non-adiabatic gravitational waveforms are one of the original aims of the self-force community, since they express i) the physics closer to the black hole horizon; ii) the most complex trajectories; iii) the most tantalising theoretical questions.  The head-on collisions of black holes and the associated radiation reaction were evoked recently in the context of particle accelerators and thereby showing the richness of the applicability of the radial trajectory also beyond the astrophysical realm. As gravity is claimed by some authors to be the dominant force in the transplanckian region, the use of general relativity is adopted for their analysis.  This chapter reviews the problem of free fall of a small mass into a large one, from the beginning of science, whatever this may mean, to the application of the self-force and of a concurring approach, in the last fourteen years. There is no pretension of exhaustiveness and, furthermore, justifiably or not, from this review some topics have been disregarded, namely: any orbit different from radial fall; radiation reaction in electromagnetism; but also the head-on of comparable masses and Kerr geometry, post-Newtonian (pN) and effective one-body (EOB) methods; quantum corrections to motion.  Herein, the terms of self-force and radiation reaction are used rather loosely, though the latter does not include non-radiative modes. Thus, the self-force describes any of the effects upon an object's motion which are proportional to its own mass. Nevertheless, to the term self-force is often associated a specific method and it is preferable to adopt the term back-action whenever such association is not meant.  Geometric units ($G = c = 1$) and the convention (-,+,+,+) are adopted, unless otherwise stated. The full metric is given by ${\\bar g}_{\\alpha\\beta} (t,r) = g_{\\alpha\\beta}(r) + h_{\\alpha\\beta} (t,r)$ where $g_{\\alpha\\beta}$ is the background metric of a black hole of mass $M$ and $h_{\\alpha\\beta}$ is the perturbation caused by a test particle of mass $m$. ",
        "conclusions": "It has been shown that free fall is still the arena for a deeper comprehension of gravitation. Furthermore, it still generates acute observations like relativistic gliding for which the asymmetric oscillations of a quasirigid body slow down or accelerate its fall in a gravitational background \\cite{gumo07}.  But is the problem of radial fall solved? Do we know the laws of motion of a star falling into a black hole, the relativistic modern version of the falling stone? A fair and objective answer leads to a moderate optimism. The general relativistic problem has had undeniable progress from 1997, but a careful analysis of the literature shows that some issues are either still partly open or simply not fully at hand, in terms of clear procedures, by means of which clearly cut answers are obtained.  The remaining steps to be fulfilled for a satisfactory level of comprehension for radial fall of a small particle into a large mass (represented by a SD black hole) within the first perturbative order in $m/M$ are divided in three groups.  Investigations to be pursued before recurring to the self-consistent prescription:  \\begin{enumerate}[I] \\item{ Compute in proper time the first order perturbation equation with deviation terms, i.e. all terms of eq. (\\ref{eq:p13}).} \\item{ Identify and compare expressions for renormalised, semi-renormalised and unrenormalised accelerations in the SD perturbed geometry; evaluate repulsive effects. The last century debate on repulsion has not addressed the effects of mass finitude ($l=0,1$) and radiation reaction on the motion ($l\\geq 2$). The inclusion of perturbations blurs the issue further and renders it more complex. } \\item{Compare the waveform corrections obtained with energy-balance by Aoudia \\cite{ao08} with those to be obtained by the integration in the equation of motion of the approximate analytic expression of the self-force of Barack and Lousto \\cite{balo02}. The comparison is of interest only in case of concurring outcomes, as both procedures have a large degree of approximation. An other comparison \\cite{sp09} could be made with the numerical results for a 10 : 1 mass ratio, as larger mass ratios imply the solution of the numerical two-scale problem, not yet at hand.} \\item{Solve the initial conditions for particles with non-null or even relativistic initial velocities with and without large eccentricities. If the outcome of the previous item is successful, coarsely identify possible signatures of radiation reaction\\footnote{For Mino and Brink \\cite{mibr08}, the energy and momentum radiated are computed on the assumption that the small body falls in a dynamical time scale, with respect to proper time, well short of the radiation reaction time scale and therefore the gravitational radiation back-action on the orbit is considered negligible. The particle plunges on a geodesic trajectory, incidentally starting from a circular orbit, thus at zero initial radial velocity.}. It is likely that a non-null initial velocity would act as an amplifier of the waveform shifts due to radiation reaction. In \\cite{lopr97a}, Lousto and Price show that the energy radiated by a non-null initial velocity may rise up to almost two orders of magnitude (see also fig. 3 in \\cite{lo97}). Incidentally, this line of work is confluent with the interests shown recently in the particle physics community \\cite{ardodsto08, spcaprbego08, shokya08, ve10, yosh09} for head-on collisions and the associated radiation reaction \\cite{gakospto10a, gakospto10b}. Hopefully, progress in numerical relativity may lead to analysis of larger mass ratios and thereby test future results of perturbation theory\\footnote{For head-on collisions, Price and Pullin have surprisingly shown the applicability of perturbation theory for the computation of the radiated energy and of the waveform for two equal black holes, starting at very small separation distances \\cite{prpu94}, the so called close-limit approximation.} in the range beyond $10^{-3}$.}  \\vskip 5pt  Investigations to be performed solely by the implementation of the self-consistent prescription:  \\vskip 5pt  \\item{Evaluate the trajectory by means of the Gralla and Wald self-consistent method \\cite{grwa08}. At each integration step, for a given number of modes: evaluation of the perturbation functions at the position of the particle; regularisation by mode-sum or $\\zeta$ methods; correction of the geodesic and identification of the cell crossed by the particle; computation of the source term; reiteration of the above. Perturbations and self-force analysis in de Donder's gauge are on-going \\cite{balo05, basa10}.} \\item{Repeat steps 1-4 on the basis of the acquired self-consistent trajectory. An iterative scheme may be envisaged also in coordinate time.}  \\vskip 5pt  Investigations to be performed independently from the self-consistent prescription:  \\vskip 5pt  \\item{Identify whether in a thought radial fall experiment, there is a physical observable independent of the gauge choice or else manifesting a recognisable effects in a given gauge, following Detweiler \\cite{de08}.} \\item{Identify the domain of applicability of the self-force.  Although self-consistency represents a closer description to perturbed motion, it is limited to cases where local deviations are small, since after all, it remains a perturbative approach. Thus, a quantitative identification of the domain of applicability of such prescription would be of interest to the community developing LISA templates. The adiabatic approximation can't be evoked to establish the limits of any self-force based analysis. Indeed, the hypothesis on the feeble magnitude of local deviations may be less constraining than the adiabatic hypothesis, for the absence of requirements on averaging. An explicit condition, referring to the geometry of the orbit and to the mass ratio and defining the domain of applicability of the self-force versus fully numerical approaches, is not yet available.} \\item{Investigate other gauges than Regge-Wheeler and other methods than the self-force for radial fall. A confirmation of the results by Barack, Lousto in \\cite{lo01, balo02} by an independent group is missing, though in \\cite{ao08} the necessary basic tools are developed, e.g. a numerical programme confirming the waveforms of \\cite{lopr97b, mapo02}. It would be also beneficial to carry out an analysis by the EOB method \\cite{da09b} \\footnote{A non-recent analysis by Simone, Poisson and Will \\cite{sipowi95}, between pN and perturbation methods in head-on collisions, was limited to the computation of the gravitational wave energy flux.}. } \\end{enumerate}  The self-force community and even more the EMRI community are animated by different interests ranging from fundamental physics and theory, to numerical applications, data analysis and astrophysics. To these variegated communities and generally to physicists and astrophysicists, the capture of stars by supermassive black holes will bring a development comparable to the ancient markings made by the leaning Pisa tower, the Cambridge apple tree and the Einstein lift.  \\begin{acknowledgement} Discussions throughout the years with S. Aoudia, L. Barack, S. Chandrasekhar, S. Detweiler, C. Lousto, J. Martin-Garcia, S. Gralla, E. Poisson, R. Price, R. Wald, B. Whiting are acknowledged. I would like to thank E. Lallier Verg\\`es for the support to the CNRS School on the Mass and the $11^{th}$ Capra conference, held in June 2008 in Orl\\'eans. It was my sincere hope that both events could put together different communities working on mass and motion: part of the contributors to this book and their colleagues (Blanchet, Detweiler, Le Tiec and Whiting) have already made concrete steps towards such cooperation \\cite{bldeltwh10a, bldeltwh10b}. \\end{acknowledgement}"
    },
    "1005/1005.5491_arXiv.txt": {
        "abstract": "This short review is addressed to cosmologists.% \\footnote{Plenary talk at the \\emph{Invisible Universe} conference held in Paris in July 2009. To appear in the Proceedings, edited by J. M. Alimi et al  (AIP Publications)}  General relativity predicts that space-time comes to an end and physics comes to a halt at the big-bang. Recent developments in loop quantum cosmology have shown that these predictions cannot be trusted. Quantum geometry effects can resolve singularities, thereby opening new vistas. Examples are: The big bang is replaced by a quantum bounce; the `horizon problem' disappears; immediately after the big bounce, there is a super-inflationary phase with its own phenomenological ramifications; and, in presence of a standard inflaton potential, initial conditions are naturally set for a long, slow roll inflation independently of what happens in the pre-big bang branch. % ",
        "introduction": "\\label{s1}  At this conference we heard of the spectacular progress that has occurred in observational cosmology in recent years. We also learned about the upcoming missions that are poised to provide new data to further constrain or even rule out leading theoretical models. These advances have been and continue to be the engines that drive contemporary cosmology. They have brought to forefront the astonishing success of the Friedmann, Lema\\^{\\i}tre, Robertson, Walker (FLRW) models, and perturbations thereof. Indeed, it appears that the rich data that we now have, and are likely to accumulate in the near future, would be adequately described by these simple applications of general relativity and quantum field theory on the resulting cosmological backgrounds.  However these theories are conceptually incomplete. They assume that the universe begins with a big bang at which matter densities and space-time curvature become infinite. With inflationary scenarios there were initial hopes that perhaps the big-bang singularity could be avoided because the inflaton fails to satisfy the strong energy condition often used in the singularity theorems of general relativity. However, Borde, Guth and Vilenkin \\cite{bgv} soon established that this hope was misplaced. Inflation can be eternal in the future but, if we go back in time using Einstein's equations, one again finds that the space-time ends and physics simply comes to a halt at the big bang. But all our experience with fundamental physics suggests that this cannot be the situation in the real world. This must be a prediction of a theory that has been pushed well beyond the domain of its validity. To know what \\emph{really} happened near the putative big bang, we must work with a genuine unification of general relativity and quantum physics, an unification which does not pre-suppose that space-time is a smooth continuum, and which can encompass the rich non-linear structure of strong gravity that lies beyond the scope of perturbation theory.  But the burden on this desired unification is heavy. In the context of cosmology there is a long list of fundamental questions that must be satisfactorily addressed by such a theory. Here is an illustrative list encompassing some of the contemporary issues. \\medskip \\begin{quote} \\noindent$\\bullet$ If general relativity is transcended, how close to the putative big bang does a smooth space-time of Einstein's make sense? In particular, can one show from first principles that this approximation is valid at the onset of inflation?\\\\ $\\bullet$ Is the big-bang singularity naturally resolved by the quantum version of Einstein's equations? Or, is some external input such as a new principle or a boundary condition at the big bang essential? An outstanding example of such an external input is the Hartle-Hawking proposal \\cite{hh}.\\\\ $\\bullet$ Is the quantum evolution across the `singularity' deterministic? One needs a fully non-perturbative framework to answer this question in the affirmative. In the pre-big-bang \\cite{pbb} and ekpyrotic/cyclic \\cite{ekp1,ekp2} scenarios, for example, so far the answer is in the negative because these theories pre-suppose the space-time continuum of general relativity and this approximation fails at the big-bang.\\\\ $\\bullet$ If the singularity is resolved, what is on the `other side'? Is there just a `quantum foam', far removed from any classical space-time, or, is there another large, classical universe? \\end{quote}  Such questions are fundamental and must be faced squarely because the resolution of classical singularities can profoundly shift the paradigm underlying contemporary cosmology. This in itself makes it imperative that we understand the quantum nature of the big bang. But there could also be another rich pay-off: the new paradigm may enable us to address open issues that are observationally significant. For example, if there is a classical pre-big-bang branch to the universe, the horizon problem would disappear and the observed large scale homogeneity could be simply a consequence of the fact that even the most distant parts of the universe would have been in causal contact in the past. If there is a pre-big-bang branch, we would not have to specify the initial conditions for perturbations on the singularity, where the applicability of current theories is least reliable. There would be more natural ways of specifying these conditions which, in turn, may well lead to small, potentially observable deviations from current predictions. Finally, our experience with general relativity itself suggests that, once the physics of the Planck regime is well understood, we may be handed with novel predictions of central importance to the next generation of astrophysicists and cosmologists.  Loop quantum gravity (LQG) \\cite{alrev,crbook,ttbook} is well suited to embark on this mission because it does not pre-suppose a classical space-time --it is background independent. At a fundamental level, everything, \\emph{including geometry,} is described in the paradigm of quantum physics \\cite{almmt,rs,al5,alvol}. Classical space-times emerge only on coarse graining of semi-classical quantum states. Finally, since the approach is fully non-perturbative, it is well suited for the strong field regime near the putative big bang.  In this chapter I will discuss loop quantum cosmology (LQC), the application of the principles of LQG to cosmology \\cite{mb-rev,aa-badhonef}. Initial ideas appeared in \\cite{mb1,abl} and were developed in detail for a variety of cosmological models in \\cite{aps1,aps2,aps3,ps,apsv,kv,acs,cs1,bp,ap,aps4,hybrid,awe2,awe3,cs2}. \\emph{In all cases, the big-bang and big-crunch singularities are resolved in a direct physical sense.} The resulting Planck scale physics has been explored using analytical and numerical solutions to the quantum Einstein equations as well as effective equations which capture the leading quantum corrections. Singularity theorems are avoided \\emph{not} because one uses matter violating energy conditions. Indeed, in the models that have been studied in most detail, \\emph{all energy conditions hold}. The theorems are inapplicable because quantum geometry effects modify Einstein equations themselves.  Physically, quantum geometry gives rise to a new repulsive force. This force is \\emph{utterly} negligible under normal circumstances. It is only when the matter density becomes about 1\\% of the Planck density or curvature approaches $1/\\lp^2$ that the repulsive force ---and hence the deviation from classical general relativity--- becomes significant. But then the repulsive force rises \\emph{very} quickly, overwhelms classical attraction and causes a quantum bounce. The density and curvature start falling and once they are below the scales just mentioned, the force again becomes negligible and classical general relativity again becomes an excellent approximation.  Immediately to the future of the bounce there is a robust phase of super-inflation which is not encountered in general relativity \\cite{mb2,ps2}. But it is short lived and in absence of a suitable inflaton potential it does not yield a sufficient number of e-foldings. However, in presence of suitable potentials ---such as $m^2\\phi^2$--- super-inflation \\emph{funnels the phase space trajectories to initial conditions which virtually guarantee a slow roll inflation with 60 or more e-foldings} \\cite{as}. This is in striking contrast to what happens in general relativity where it has been argued \\cite{gt} that the probability of $N$ e-folding decreases as $e^{-3N}$. The super-inflationary phase is also likely to have other phenomenological consequences ---such as production of gravitational waves--- that are being analyzed.  The article is organized as follows. Section \\ref{s2} lays out the conceptual setting and section \\ref{s3} provides a bird's eye view of LQC through illustrative results. We will see that not only has LQC answered many of the long standing questions but it has also opened new vistas. Section \\ref{s4} summarizes the origin of the novel predictions and places them in a broader perspective. ",
        "conclusions": ""
    },
    "1005/1005.2278_arXiv.txt": {
        "abstract": "{} {We test predictions of evolution models on mixing of CNO-cycled products in massive stars from a fundamental perspective. Relative changes within the theoretical C:N:O abundance ratios and the buildup of helium are compared with observational results.} {A sample of well-studied Galactic massive stars is presented. High-quality optical spectra are carefully analysed using improved NLTE line-formation and comprehensive analysis strategies. The results are put in the context of the existing literature data.} {A tight trend in the observed $N/C$ vs. $N/O$ ratios and the buildup of helium is found from the self-consistent analysis of main-sequence to supergiant stars for the first time. The catalytic nature of the CNO-cycles is confirmed quantitatively, though further investigations are required to derive a fully consistent picture. Our observational results support the case of strong mixing, as predicted e.g. by evolution models that consider magnetic fields or by models that have gone through the first dredge-up in the case of many supergiants.} {} ",
        "introduction": "Energy production in massive stars is governed by the CNO-cycles throughout most of their lifetimes. The general correctness of our understanding of the CNO cycles and of the relevant nuclear data \\citep[e.g.][]{maeder83} is confirmed impressively by observation: when massive stars enter the Wolf-Rayet phase as WN subtypes, equilibrium CNO-processed material becomes exposed on their surface \\citep[see e.g.][]{crowther07}.  However, traces of mixing of CNO-cycled products from the stellar core to the stellar surface can already be found much earlier in the lives of massive stars. Observational indications of superficial abundance anomalies for carbon, nitrogen, and oxygen (and the burning product helium) in OB-type stars even on the main sequence (MS) and, more prominently, in the blue supergiants was found early in classification spectrograms \\citep[e.g.][]{walborn76}. Subsequent analyses provided evidence for a characteristic enrichment of nitrogen -- which is the easiest to be detected -- and helium in many early-type stars, both near the MS and in blue supergiants \\citep[e.g.][]{schoenberner88,gila92,herrero92,kilian92,venn95,lyubimkov96,mcerlean99}.  A theoretical understanding of {\\em early} mixing of CNO-pro\\-cessed material to the stellar surface could not be achieved within the framework of evolution calculations for non-rotating stars with mass-loss, which were state-of-the-art at that time \\citep[e.g.][]{chma86}. The pollution of the surface layers with CNO-cycled material only occurs when the star reaches the red supergiant phase in such models, via convective dredge--up.  Considered a secondary effect for a long time, stellar rotation has come lately into focus. It turned out that rotationally-induced mixing through meridional circulation and turbulent diffusion in rotating stars provides the means to change all model outputs substantially and to bring theory and observation into much better agreement \\citep{mame00,hela00}. The latest step taken in the modelling was to consider the effects from an interplay of rotation and magnetic fields, which -- depending on the detailed input physics and approximations made -- only provide minor modifications to the surface abundances \\citep{heger05} or substantial changes \\citep[henceforth abbreviated as MM05]{mame05}. These differences result mainly from the two groups using different sets of equations for computing the effects of magnetic fields; see in particular, the changes brought to the system of equations given by \\citet{spruit02} by MM05 in their Sect.~2.  The only means to verify the models is via systematic comparison with observations covering the relevant parts of the Hertzsprung-Russell diagram. Homogeneous analyses of a larger star sample from the main sequence to the supergiant stage have only recently become available. The results in particular for N abundances apparently challenge the concept of rotational mixing in massive stars \\citep{hunter09} and thus the present-day evolution models \\citep[see, however,][]{maederetal09}.  \\begin{figure*} \\centering \\hfill \\includegraphics[width=.485\\textwidth]{14164fg1a.eps}\\hfill \\includegraphics[width=.485\\textwidth]{14164fg1b.eps} \\hfill \\caption{Status of observational constraints on the (magneto-)hydrodynamic mixing of CNO-burning products in massive stars from previous NLTE analyses. Mass ratios $N/C$ over $N/O$ are displayed. Left panel: main-sequence stars. Circles: \\citet{kilian92}; triangles: \\citet{gila92}; diamonds: \\citet{cula94}, \\citet{daflon99,daflon01a,daflon01b}; squares: \\citet{morel08}; crosses: \\citet{hunter09}. Right panel: BA supergiants. Triangles: \\citet{venn95}, \\citet{vepr03}; circles: \\citet{takeda00}; squares: \\citet{crowther06}; diamonds: \\citet{searle08}. Error bars are omitted for clarity: uncertainties in the abundances of the individual elements are typically about a factor 2, such that the error bars can be larger than the plotting range. The lines represent predictions from evolution calculations, for a rotating 15\\,$M_\\odot$ star \\citep[$v_\\mathrm{rot}^\\mathrm{ini}$\\,=\\,300\\,km\\,s$^{-1}$,][MM03; until the end of the MS: solid red line, until the end of He burning: dashed blue line]{mema03} and for a star of the same mass and $v_\\mathrm{rot}^\\mathrm{ini}$ that in addition takes the interaction of rotation and a magnetic dynamo into account \\citep[MM05; until the end of the MS: dotted line]{mame05}, respectively. The predicted trends are similar for the entire mass range under investigation, see Fig.\\ref{cnomix920}.} \\label{litsummary} \\end{figure*}  Here, we address the topic of early mixing of CNO-cycled products in massive stars from a fundamental perspective. We compare the model predictions for the relative changes within the C:N:O abundance ratios and the buildup of helium (Sect.~\\ref{theory}) with observations. For this we briefly review the status of the literature on CNO abundances and introduce a well-studied Galactic star sample for which high-accuracy analyses of high-quality spectra were performed using improved NLTE modelling and comprehensive analysis techniques (Sect.~\\ref{obs}). Conclusions from this comparison are drawn in Sect.~\\ref{conclusions}. ",
        "conclusions": "} Evolution models for massive stars predict very tight relations for the change of surface $N/O$ and $N/C$ abundance ratios and for the buildup of helium, as a consequence of mixing with CNO-cycled matter (for given initial chemical composition). Massive stars in the solar neighbourhood, which supposedly share a nearly uniform initial composition (PNB08), are therefore expected to follow the predicted relations, which are governed by nuclear reactions and the dilution effects produced by mixing.  A comparison with NLTE spectral analyses from the literature on CNO abundances in massive stars of $\\sim$8--25\\,$M_\\odot$ leaves room for broad interpretation because of large uncertainties. Observations may even be interpreted as posing a challenge to theory, in particular when the full set of information is not accounted for and conclusions are drawn from one indicator alone, e.g. nitrogen.  On the other hand, our high-precision analyses of a sample of Galactic massive stars from the main sequence to the supergiant stage find these tight correlations. Even though further investigations are required to refine the observational constraints on helium abundances before a fully coherent picture can be obtained, the case for strong mixing is clearly supported. It is predicted, e.g., by the models of MM00, by models with magnetic field (MM05), or in cases where the models have gone through the first dredge-up. The separation of the different possibilities may come from studing the evolution of the rotation velocities with time. On the one hand, models with magnetic field predict higher rotational velocities at the end of the MS phase, because the strong internal coupling transmits some of the fast core rotation to the surface. On the other hand, stars on a blue loop should be slower rotators than objects on their first passage from blue to red on average because of the additional angular-momentum loss experienced through strong mass loss during the red supergiant phase. Indeed, some of the slowest rotators show very strong mixing signatures (Table~\\ref{table1}).  Finally, we have to point out that some of our results have the potential of challenging the currently available evolution models. The star $\\tau$\\,Sco (HD149438) stands out in the sample as it shows characteristics that may be explained by a homogeneous evolution, but it requires a highly-efficient spin-down mechanism. One may speculate on magnetic breaking due to angular-momentum losses by a magnetically confined line-driven stellar wind or magnetic coupling to the accretion disc during the star-forming process in the case of a fossil field. Even though the topic is not understood theoretically in a comprehensive way, spin-down times of the order of 1\\,Myr \\citep{ud-Doula09} or even less \\citep{mikulasek08} are reported for some magnetic massive stars, possibly leading to the required slow rotation already on the zero-age MS in this case as well. Then, if the present helium abundances are confirmed and the supergiants are shown to have evolved directly from the MS, a different kind of mixing may also be required.  We conclude that the tight observational constraints that are required for a thorough testing of the stellar evolution models are within reach. Applications of the improved modelling and analysis techniques to high-quality observations of larger star samples in the Milky Way and other galaxies may finally provide the empirical basis to benchmark the models. It may thus become feasible to disentangle the effects of metallicity, rotation, magnetic fields, and binarity on massive star evolution."
    },
    "1005/1005.3754_arXiv.txt": {
        "abstract": "{Oxygen is the third most common element in the Universe. The measurement of oxygen lines in metal-poor unevolved stars, in particular near-UV OH lines, can provide invaluable information about the properties of the Early Galaxy.} {Near-UV OH lines constitute an important tool to derive oxygen abundances in metal-poor dwarf stars. Therefore, it is important to correctly model the line formation of OH lines, especially in metal-poor stars, where 3D hydrodynamical models commonly predict cooler temperatures than plane-parallel hydrostatic models in the upper photosphere.} {We have made use of a grid of 52 3D hydrodynamical model atmospheres for dwarf stars computed with the code \\cobold, extracted from the more extended CIFIST grid. The 52 models cover the effective temperature range 5000--6500\\,K, the surface gravity range 3.5--4.5 and the metallicity range $-3 < [{\\rm Fe}/{\\rm H}] < 0$.} {We determine 3D-LTE abundance corrections in all 52 3D models for several OH lines and \\ion{Fe}{i} lines of different excitation potentials. These 3D-LTE corrections are generally negative and reach values of roughly -1~dex (for the OH 3167 with excitation potential of approximately 1~eV) for the higher temperatures and surface gravities.} {We apply these 3D-LTE corrections to the individual O abundances derived from OH lines of a sample the metal-poor dwarf stars reported in Israelian et al.(1998, 2001) and Boesgaard et al.(1999) by interpolating the stellar parameters of the dwarfs in the grid of 3D-LTE corrections. The new 3D-LTE [O/Fe] ratio still keeps a similar trend as the 1D-LTE, i.e, increasing towards lower [Fe/H] values. We applied 1D-NLTE corrections to 3D \\ion{Fe}{i} abundances and still see an increasing [O/Fe] ratio towards lower metallicites. However, the Galactic [O/Fe] ratio must be revisited once 3D-NLTE corrections become available for OH and Fe lines for a grid of 3D hydrodynamical model atmospheres. } ",
        "introduction": "\\label{intro}  The metal-poor stars of the Galactic halo provide the fossil record of the early Galaxy's composition. Dwarf halo stars are particularly relevant, because their atmospheres are not significantly altered by internal mixing and provide a unique tracer to constrain Galactic chemical evolutionary models. Oxygen is a key element in this scenario, because it is the most abundant element in stars after H and He. It is produced in the interiors of massive stars by hydrostatic burning and its content is modified during the explosive nucleosynthesis in type~II supernovae (SNe) and hypernovae (HNe) and returned into the interstellar medium. On the other hand, iron is created by both type~II and type~I SN explosions. However, type~I SNe progenitors have longer lifetimes, which is why the abundance ratio of [O/\\ion{Fe}{i}] can be used to constrain the chemical evolution of the Galaxy.  There have been numerous studies of the oxygen abundance in halo stars. Despite considerable observational and theoretical efforts, the trend of [O/Fe] ratio\\footnote{$[{\\rm O}/{\\rm Fe}]=\\log [N({\\rm O})/N({\\rm Fe})]_\\star-\\log [N({\\rm O})/N({\\rm Fe})]_\\odot$} versus [Fe/H] is still unclear. The analysis of the forbidden line \\ion{O}{i} 6300~{\\AA} in giants shows a plateau with [O/Fe$] \\sim 0.4-0.5$ in the metallicity range $-2.5 < [{\\rm Fe}/{\\rm H}] < -1$ \\citep{bar88} and [O/Fe$] \\sim 0.7$ for $-4. < [{\\rm Fe}/{\\rm H}] < -2.5$ \\citep{cay04}. A similar behaviour is seen for metal-poor subgiant stars in the range $-3. < [{\\rm Fe}/{\\rm H}] < -1.5$ \\citep[{[O/Fe]} $\\sim 0.4-0.5$,][]{gar06}. However, \\citet{gar06} already noted that by plotting all measurements from the [OI] line for dwarfs \\citep{nis02}, subgiants \\citep{gar06} and giants \\citep{cay04}, the picture changes and an increasing trend [O/Fe] towards lower metallicities clearly appears.  The near infrared (IR) triplet \\ion{O}{i} 7771--5~{\\AA} in metal-poor dwarfs and subgiants \\citep{abi89,isr01,gar06} points towards increasing [O/Fe] values with decreasing [Fe/H], although the O abundances derived from the near-IR triplet are typically $\\sim 0.4-0.7$ higher than those derived from the forbidden \\ion{O}{i} line \\citep{ful03}. The OH A-X electronic lines in the near ultraviolet (UV) provide also higher [O/Fe] ratios towards lower [Fe/H] values in dwarf stars \\citep{isr98,boe99,isr01,gon08}. However, \\citet{gar06} found a quasi-plateau of [O/Fe] for subgiant stars in the range $-3. < [{\\rm Fe}/{\\rm H}] < -1.5$, using Fe abundances determined from \\ion{Fe}{ii} lines. The [O/Fe] ratio shows a negative slope if one instead uses the Fe abundances estimated from \\ion{Fe}{i} lines, although with lower [O/Fe] values than those determined for metal-poor dwarfs. It is advisable to use \\ion{Fe}{i} lines instead of \\ion{Fe}{ii} lines to derive the [O/Fe] ratio because of its similar sensitivity to the surface gravity.  These O-abundance indicators present different complications. The near-IR \\ion{O}{i} triplet is susceptible to non-local thermodynamical equilibrium (NLTE) effects \\citep[and references therein]{kis01}, with abundance corrections below 0.2 dex, and is quite sensitive to the adopted \\teffo. The [OI] is not sensitive to departures from LTE, but it is essentially undetectable in dwarfs with $[{\\rm Fe}/{\\rm H}]\\lesssim -2$. The near-UV OH lines are strongly sensitive to the temperature structure and inhomoginities \\citep{asp01,gon08}. In addition, \\ion{Fe}{i} lines suffer from severe NLTE effects in metal-poor stars \\citep[see e.g.][]{the99}. We note that \\citet{shc05} have performed NLTE computations for the metal-poor subgiant HD140283 ($[{\\rm Fe}/{\\rm H}]\\sim -2.5$) with a single snapshot of a 3D hydrodynamical simulation \\citep{asp99}, and found NLTE-LTE corrections of $+0.9$ and $+0.4$ for \\ion{Fe}{i} and \\ion{Fe}{ii} lines, respectively.  Finally, the oxygen abundance in the Sun is still a matter of debate. We will adopt throughout this work the value of $\\log [N({\\rm O})/N({\\rm H}]_\\odot =8.76$, which was determined through 3-dimensional (3D) hydrodynamical models \\citep{caf08}.  We have used a subset of the CIFIST grid of 3D hydrodynamical model atmospheres \\citep{lud09} to investigate the 3D-LTE and 3D-NLTE [O/Fe] and [O/H] trends in metal-poor dwarf stars. ",
        "conclusions": "The large grid of 3D hydrodynamical model atmospheres of dwarf stars quite captivated our attention during the last three years \\citep{lud09}. It took thousands of hours in computing time to build such a grid. We have used 52 3D models extracted from this grid with the following stellar parameters and metallicities: \\teff$=5000$, 5500, 5900, 6300 and 6500~K, \\logg$=3.5$, 4 and 4.5, and [Fe/H$]=0$, --1, --2 and --3. The main difference with the ``classical'' 1D models is that 3D models show temperature inhomogeneities and a cooler average temperature profile in models with [Fe/H$]<-1$.  This allowed us for the first time to compute 3D abundance corrections of several near-UV OH and \\ion{Fe}{i} lines. These 3D corrections are generally larger for higher effective temperatures, larger surface gravities, and lower metallicities. In addition, lines with lower excitation potentials show stronger 3D corrections.  We applied this grid of 3D corrections to a sample of metal-poor dwarf stars from \\citet{isr98,isr01}, \\citet{boe99} and the most metal-poor dwarf stars of the binary \\mbox{CS 22876--032} from \\citet{gon08}. We interpolated within this grid, using the stellar parameters and metallicities of these stars.  Finally, we are able to display a new trend with the 3D-NLTE [O/\\ion{Fe}{i}] ratio versus metallicity and this trend still increases towards lower metallicities, as the 1D-NLTE, but with a smaller slope in absolute value. However, we caution that some assumptions (as e.g. the restriction to a 6-bin scheme in most 3D models and NLTE corrections only in 1D), which we believe to be reasonable, have been adopted to achieve this result and that therefore, this result should be taken with that in mind. A full 3D-NLTE study of all 3D model atmospheres presented in this paper must be performed for both near-UV OH lines and \\ion{Fe}{i} lines to see if this trend changes."
    },
    "1005/1005.4668_arXiv.txt": {
        "abstract": "The PAMELA satellite has observed an excess of positrons over electrons in the energy range $1-100~GeV$ that increases with energy. We propose that the excess is not due to a change in the local interstellar spectrum, but is due to heliospheric modulation. We motivate this from the known form of the heliospheric magnetic field and predict that the excess will disappear when we enter a period of solar maximum activity. ",
        "introduction": "The PAMELA satellite \\cite{Adriani:2009p431} has observed a rise in the positron to electron fraction that increases with energy above $10~GeV$. This has been interpreted as proof of a rise in the positron fraction in the local interstellar spectrum of cosmic rays. Such a result requires that there be some hardening of the positron spectrum at these energies and has caused a great deal of discussion of possible sources from decaying or annihilating dark matter to acceleration from a population of pulsars.  We propose instead that this rising fraction is due to the heliospheric magnetic field. The heliospheric magnetic field is ordered over large distances during solar minimum \\cite{Zhou:2009p1228,UlyssesBook} and is well described up to medium latitudes by a Parker spiral field \\cite{Parker:1958p1139}. During solar minimum this means that for a particle to propagate directly to earth it must traverse a large scale ordered magnetic field that is primarily oriented perpendicularly to the radial. Direct transport is impossible for low energy particles ($E<1~GeV$) as their Larmor radius is small even in the outer heliosphere. They instead reach the central solar system primarily by following the magnetic field lines along with a complex process involving convection, diffusion, drifts and energy loss. At high energies ($E>1~TeV$), particles can traverse the ordered magnetic field with little deviation. At these energies we see the interstellar spectrum directly.  If the magnetic field were the only effect on the particle transport, and if the magnetic field were time invariant, then the switch from one set of particles paths to the other would have no effect on the flux by Liouville's theorem. However the magnetic field is time variant. The magnetic field lines are carried by the supersonic solar wind and travel out towards the edge of the heliosphere at 400~km/s. This causes large scale convection of cosmic rays that results in a deficit in the cosmic ray density in the central solar system when compared to the interstellar spectrum. This has been well measured for cosmic rays with rigidities up to 2.5GV \\cite{UlyssesBook,Ndiitwani:2005p2956}. At these rigidities, the particles generally have a very long path length to the center of the solar system as they follow the magnetic field lines. Therefore they have a low radial velocity, and are greatly effected by the convection of the solar wind. The solar magnetic field also varies on longer timescales, affected by the 27 day rotation period of the sun, and by the reversal of the solar magnetic field every 11 years.  At higher rigidities particles can travel directly to the central solar system at the speed of light. Their path length within the heliosphere is short and they are almost unaffected by convection from the solar wind. Therefore at high energies we expect there to be no radial gradient of cosmic rays as they should free stream through the solar system. From these two observations - the two separate transport paths for low and high energy particles, and radial gradient of low rigidity cosmic rays from the convection of the solar wind - we can already say that there must be a hardening of the cosmic ray spectrum between rigidities of $1~GV$ and $1000~GV$. This will be due to the change over from a locally suppressed cosmic ray flux to the interstellar flux.  To explain the PAMELA signal we need one further element. We require that the change over from the local flux to the interstellar flux be charge asymmetric. We already know that the orientation of the magnetic field results in asymmetric cosmic ray transport. Indeed this is the cause of the low energy heliospheric modulation below a few $GeV$\\cite{Potgieter:2001p1465}. We argue that this is also true for high energy particles and that the asymmetry should become more pronounced with increasing energy. We use a simple model of the heliospheric magnetic field to show this in section \\ref{lense1}. This exhibits the primary features of the asymmetry that we are interested in. We then extend the study to consider the asymmetry in a more complex realisation of the field in section \\ref{lense2}.  By postulating a heliospheric cause of the positron excess we can immediately make a clear prediction. We predict that the excess should disappear when the solar system enters a period of solar maximum activity which it is due to do shortly. In addition, in the opposite solar cycle when the sun's magnetic field is oppositely aligned, we would expect to see the opposite modulation - giving a rising excess of electrons over positrons in the energy range $1-100~GeV$. We will also show that the tilt angle of the heliospheric current sheet (HCS) is an important feature of the model, and predict that any temporal variation of the latitudinal extent of the current sheet should be matched by a variation in the observed positron excess, with a time delay due to the time taken for magnetic field features at the sun to propagate out to the heliopause in the solar wind. These predictions are robust features of our model and testable in the near future. ",
        "conclusions": "}  We have proposed that the PAMELA positron excess is not down to annihilating dark matter or new astrophysical sources, but instead results from the configuration of the heliospheric magnetic field. The ordered nature of the magnetic field on large scales creates a lense that allows particles of one sign to free stream into the center of the solar system whilst particles off the opposite sign travel out. This effect rises with energy. It naturally occurs in the correct energy range for the correct sign of particle to account for the observed positron excess. Though our model of the solar system magnetic field is rough, we can already make two clear predictions. Firstly, the PAMELA result will disappear when we go through solar maximum (with a suitable delay to account for the time it takes the conditions at the sun to propagate out through the solar system). Secondly, we predict that there will be fluctuations in the observed excess that correlate (again with a delay) to the tilt angle of the heliospheric current sheet."
    },
    "1005/1005.4497_arXiv.txt": {
        "abstract": "We study mass transfer by Roche lobe overflow in close-in exoplanetary systems. The planet's atmospheric gas passes through the inner Lagrangian point and flows along a narrow stream, accelerating to 100-$200~\\kms$ velocity before forming an accretion disk. We show that the cylinder-shaped accretion stream can have an area (projected in the plane of the sky) comparable to that of the planet and a significant optical depth to spectral line absorption. Such a ``transiting cylinder'' may produce an earlier ingress of the planet transit, as suggested by recent HST observations of the WASP-12 system. The asymmetric disk produced by the accretion stream may also lead to time-dependent obscuration of the starlight and apparent earlier ingress. We also consider the interaction of the stellar wind with the planetary magnetosphere. Since the wind speed is subsonic/sub-Alfvenic and comparable to the orbital velocity of the planet, the head of the magnetopause lies eastward relative to the substellar line (the line joining the planet and the star).  The gas around the magnetopause may, if sufficiently compressed, give rise to asymmetric ingress/egress during the planet transit, although more works are needed to evaluate this possibility. ",
        "introduction": "The close-in exoplanets (with period less than a few days) discovered in radial velocity and transit surveys are of great interest, not only because they constrain theories of planet formation and evolution, but also because they provide a probe of various physical processes that are otherwise unimportant in ``normal'' planets.  WASP-12b is a transiting exoplanet orbiting extremely close to a late-F/early-G star ($M_\\star=1.35M_\\sun$, $R_\\star=1.57R_\\sun$, $T_{\\rm eff}=6300$~K), with the orbital period $P=1.09$~days and orbital semi-major axis $a=0.023~{\\rm AU}=4.94R_\\sun =3.15R_\\star$.  The planet mass $M_p=1.41M_J$ and radius $R_p=1.79R_J$, as determined by transit observation in optical continuum (Hebb et al.~2009; Campo et al.~2010). The planet is one of the most irradiated exoplanets (with equilibrium temperature $T_{\\rm eq}=2500$-3000~K) and is highly inflated. The fact that the Roche radius (Hill sphere radius) of the planet, $R_L=a(M_p/3M_\\star)^{1/3}=1.85R_p$, is only slightly larger than the $R_p$ derived from optical transit measurements, suggests that mass loss from the planet is likely (Li et al.~2010). The small orbital separation also suggests that stellar wind may influence the atmosphere the planet.  Recently, Fossati et al.~(2010) obtained near-UV transmission spectroscopy of WASP-12b with the Cosmic Origins Spectrograph on the {\\it Hubble Space Telescope}. The data revealed enhanced transit depths (by about a factor of 2) in two wavelength bands, NUVA (2539-2580~$\\AA$) and NUVC (2770-2811~$\\AA$), which were attributed to flux attenuation by absorption lines of metals in the vicinity of the planet. Most interestingly, the NUVA data exhibits an earlier ingress compared to the transit in optical J,B and Z bands, while the egress of the transit occurs at about the same time as the optical transit.  The asymmetric behavior of the ingress/egress in the NUVA band relative to the continuum is difficult to understand if the absorbing gas surrounding the planet arises entirely from an irradiation-driven wind -- such a wind has been studied extensively in the context of hot Jupiter HD 209458b (e.g., Yelle 2004; Tian et al.~2005; Garcia Munoz 2007; Murray-Clay et al.~2009): The wind is most strongly generated at the planet's dayside, and would be distributed on both sides (terminations) of the substellar line (the line joining the planet and the star). If anything, it would be preferentially on the west side because of the planetary rotation, which is almost likely synchronized with the orbit (see Schneiter et al.~2007 for a simulation).  In this paper, we consider two possible explanations for the asymmetric excess absorption during WASP-12b transit. First, we study the Roche lobe overflow from the planet to the parent star (see Li et al.~2010) and the associated accretion stream. The stream is asymmetric with respect to the line joining the planet and the star. We show that the stream has a sky-projected area comparable to the projected planet area, and a sufficiently large column density to cause blockage of the star light prior to optical ingress. Second, we qualitatively discuss the magnetopause produced by the interaction between the stellar wind and the planet's magnetosphere. Because the planet's orbital velocity, $v_{\\rm orb}=(GM_\\star/a)^{1/2} =228~\\kms$, is comparable to the stellar wind velocity ($v_w\\simeq 100~{\\rm km~s}^{-1}$ at $a=0.023$~AU), the head of the magnetopause lies eastward relative to the substellar point. Again, ingress/egress asymmetry may be produced if the gas density in around the magnetopause is sufficiently high.   This paper is organized as follows. In Sect.~2 we discuss line absorption by a moving medium with a velocity gradient and derive the observational constraints on the absorbing gas in WASP-12b. In Sect.~3 we study the property of the accretion stream and the obscuration of the star light by the stream. Section 4 examines the possibility that the gas around the magnetopause may absorb the star light. We conclude in Sect.~5. ",
        "conclusions": "Motivated by the recent near-UV spectral observations of the WASP-12 planetary system (Fossati et al.~2010), which showed an earlier ingress of the planet's transit compared to the optical continuum, we have studied two possible explanations for the ingress/egress asymmetry. The first involves mass transfer through Roche lobe overflow, which results in an elongated accretion stream in which gas flows from the L1 Lagrangian point toward the star. Our analysis of the geometric and velocity structure of the stream shows that it may indeed provide asymmetric obscuration of the star light during the planet transit. In addition, it is also possible that an asymmetry in the accretion disk, caused by the impact of the accretion stream, produces an apparent earlier ingress. Another possibility involves the magnetopause separating the stellar wind and the planetary magnetosphere. Because of the planet's orbital motion, the head magnetopause lies eastward relative to the substellar point. We suggest that line absorption by the gas around the magnetopause may also explain the asymmetric ingress/egress behavior, although more theoretical works are needed to understand the property of the absorbing gas around the magnetopause.  The two possibilities studied in this paper may be distinguished by the fact that in the accretion stream, the gas falls away from the observers, producing redshifted absorption, while the flow around the magnetosphere tends to move toward the observer, thus producing blueshifted absorption.  Finally, we note the observed behaviors reported by Fossati et al are only marginally significant (at most $3\\sigma$ effects). Thus, more observations will be useful. Nevertheless, the physical processes discussed in this paper are quite general and therefore may be applicable in other close-in exoplanetary systems.       \\medskip We thank the participants of the morning coffee of the KITP Exoplanets program (2010.1-2010.5) for lively discussions. D.L. especially thanks Phil Arras, Doug Lin and Steve Lubow for useful discussion during the final phase of our work, and P. Arras for his comments on an earlier draft of our paper. C.H. thanks Luca Fossati for discussion of early-ingress observations. Part of this work was performed while the authors were in residence at KITP, funded by the NSF Grant PHY05-51164."
    },
    "1005/1005.3048_arXiv.txt": {
        "abstract": "We present molecular line observations, made with angular resolutions of $\\sim$20\\arcsec, toward the filamentary infrared dark cloud G34.43+0.24 using the APEX [\\dco , \\tco, \\cdo\\ and \\cssiete\\ transitions], Nobeyama 45 m [\\csdos, \\sio, \\cstyc, \\hco , \\hcotrece\\ and \\metanol\\ transitions] , and SEST [\\csdos\\ and \\cdodos\\ transitions] telescopes. We find that the spatial distribution of the molecular emission is similar to that of the dust continuum emission observed with 11\\arcsec\\ resolution (Rathborne et al. 2005) showing a filamentary structure and four cores. The cores have local thermodynamic equilibrium masses ranging from $3.3\\times10^2$ -- $1.5\\times10^3$ \\Msun\\ and virial masses from $1.1\\times10^3$ -- $1.5\\times10^3$ \\Msun, molecular hydrogen densities between $1.8\\times10^4$ and $3.9\\times10^5$ cm$^{-3}$, and column densities $>2.0\\times10^{22}$ cm$^{-2}$; values characteristics of massive star forming cores. The \\tco\\ profile observed toward the most massive core reveals a blue profile indicating that the core is undergoing large--scale inward motion with an average infall velocity of 1.3 \\kms\\ and a mass infall rate of $1.8\\times10^{-3}$ \\Msun\\ yr$^{-1}$. We report the discovery of a molecular outflow toward the northernmost core thought to be in a very early stage of evolution. We also detect the presence of high velocity gas toward each of the other three cores, giving support to the hypothesis that the excess 4.5 $\\mu$m emission (``green fuzzies'') detected toward these cores is due to shocked gas. The molecular outflows are massive and energetic, with masses ranging from 25 -- 80 \\Msun, momentum 2.3 -- $6.9\\times10^2$ \\Msun\\ \\kms, and kinetic energies 1.1 -- $3.6\\times 10^3$ \\Msun\\ km$^2$ s$^{-2}$; indicating that they are driven by luminous, high-mass young stellar objects. ",
        "introduction": "Infrared dark clouds (IRDCs) are cold ($<$25 K), massive ($\\sim$$10^2$ -- $10^4$ \\Msun), and dense ($>$$10^5$ cm$^{-3}$) molecular clouds with high column densities ($\\sim$$10^{23}$--$10^{25}$ cm$^{-2}$) seen as a dark silhouette against the bright mid--infrared background emission  (Perault et al. 1996; Egan et al. 1998; Carey et al. 1998, 2000; Hennebelle et al. 2001; Simon et al. 2006a, 2006b). A wealth of recent observations show that IRDCs are the sites of high-mass stars and star cluster formation (Rathborne et al. 2006, 2007; Pillai et al. 2006; Jackson et al. 2008; Chambers et al. 2009). Dust continuum observations reveal that IRDCs harbor compact cold cores with typical masses of $\\sim$120 \\Msun\\ and sizes $<$0.5 pc (Rathborne et al. 2006). The masses and sizes of the IRDC cores are similar to those of the massive hot cores (e.g., Garay \\& Lizano 1999), but IRDC cores are much colder. Recently, Chambers et al. (2009) investigated a sample of 190 cores found toward 38 IRDCs by Rathborne et al. (2006), and classified them as ``active'' if they are associated with enhanced 4.5 $\\mu$m emission, the so-called green fuzzies, and  an embedded 24 $\\mu$m source; and ``quiescent'' if they contain neither of these indicators. They found 37 active cores and 69 quiescent cores. Further, active cores are usually associated with bright ($>$1 Jy) methanol and water masers (Pillai et al 2006; Wang et al. 2006; Chambers et al. 2009), and have high bolometric luminosities, indicating that they are forming high--mass stars (M $>$ 8 \\Msun) stars. The quiescent cores are the most likely candidates for high--mass starless cores.  G34.43+0.24 is an IRDC with a filamentary morphology, extending by $\\sim$9\\arcmin\\ from north to south in equatorial projection (9.7 pc at the source distance of 3.7 kpc; Fa\\'undez et al. 2004; Simon et al. 2006b). This IRDC is in the list of Rathborne et al. (2006) and Chambers et al. (2009). It is located roughly 11\\arcmin\\ north of the  ultra-compact (UC) H II complex G34.26+0.15 (Molinari et al. 1996). Molecular line observations toward this cloud were first reported by Miralles et al. (1994), who from ammonia observations with 1.5\\arcmin\\ resolution found an elongated structure in the N-S direction with a total mass of 1000 \\Msun. Near the center of the filament lies the IRAS point source 18507+0121 which has a luminosity of $3.4\\times10^4$ \\Lsun.  Bronfman et al. (1996) detected toward IRAS 18507+0121 strong \\csdos\\ emission with broad line wings, indicating it is associated with a dense massive star forming region. Ramesh et al. (1997) observed the IRAS source in the \\hco, \\hcotrece, \\csdos\\ and \\cstyc\\  molecular lines, with angular resolutions of $\\sim$16\\arcsec, modeling the observed line profiles as due to emission from a collapsing hot core that is hidden behind a cold and dense envelope. From observations of 3 mm continuum and H$^{13}$CO$^+$ line emissions, with $\\sim$5\\arcsec\\ angular resolution, Shepherd et al. (2004) detected two compact molecular cores, separated by 40\\arcsec, toward the center of the filament. Millimeter continuum observations revealed the presence of four dust cores (Fa\\'undez et al. 2004; Garay et al. 2004; Rathborne et al. 2005), labeled by the latter authors as MM1, MM2, MM3, and MM4. The MM1 dust core, which corresponds to the northern compact molecular core of Shepherd et al. (2004), contains a deeply embedded luminous object surrounded by several hundred solar masses of warm gas and dust. Based on the weak 6 cm continuum emission and the lack of detection at NIR wavelenghts, Shepherd et al. (2004) suggested that the embedded object appears to be a massive B2 protostar in an early stage of evolution. The MM2 dust core, which corresponds to the southern compact molecular core of Shepherd et al. (2004), is associated with IRAS 18507+0121; with a NIR cluster of young stars with a central B0.5 star (Shepherd et al. 2004); with an UC H II region (Miralles et al. 1994; Molinari et al. 1998); with a variable H$_2$O maser (Miralles et al. 1994); and with CH$_3$OH maser emission (Szymczak et al. 2000). The MM3 dust core is located close to the northern edge of the filament, $\\sim$3.5\\arcminspace north of MM2. Garay et al. (2004) suggested that MM3 corresponds to a massive and dense cold core that will eventually collapse to form a high-mass star. The MM4 core is located $\\sim$30\\arcsecspace south of MM2 (Rathborne et al.  2005). Table~\\ref{tbl-1.2mm} lists the measured 1.2 mm dust continuum emission properties as measured by Rathborne et al. (2006). Rathborne et al. (2005) and Chambers et al (2009) reported an excess 4.5 $\\mu$m emission toward MM1, MM3, and MM4 cores, which could be produced by either ionized gas and/or shocked gas and acts therefore as a signature of current star formation (MM2 is classified by the last authors as a ``red'' core, namely with bright 8 $\\mu$m emission). They also showed that the four MM cores are associated with 24 $\\mu$m point sources (emission from warm dust), which is an indicator of accretion onto an embedded protostar. Water maser emission, a well--known tracer of low--mass and high--mass star formation, has been detected toward all four MM cores (Wang et al. 2006; Chambers et al. 2009). Chambers et al. (2009) detected Class I CH$_3$OH masers, another well--know signpost of star formation toward MM1, MM2, and MM3 cores. Rathborne et al. (2006) identified nine 1.2 mm continuum cores with masses ranging from 80 to 1300 \\Msun\\ toward G34.43+0.24.  Shepherd et al. (2007) observed the central region of the filament, encompassing the MM1 and MM2 cores, at millimeter wavelengths, in CO$(1\\rightarrow0)$, $^{13}$CO$(1\\rightarrow0)$, and C$^{18}$O$(1\\rightarrow0)$ at $\\sim4$\\arcsec\\ angular resolution and at near infrared wavelengths with the {\\it Spitzer} Space Telescope. They discovered five massive outflows, three of them associated with the MM2 core and the remaining two associated with the MM1 core. From the {\\it Spitzer} data they identified 31 young stellar objects in the filament, with a combined mass of $\\sim$127 \\Msun, and additional 22 sources that they believe could be cluster members based on the presence of strong 24 $\\mu$m emission. Cortes et al. (2008) observed MM1 and MM2 in polarized thermal dust emission at 3 mm and CO$(1\\rightarrow0)$ line emission with BIMA array. Their results suggest that there is a magnetic field othogonal to the axis formed by the MM1 and MM2 cores. Recently, Rathborne et al. (2008) carried out high angular submillimeter observations toward the MM1 core using the Submillimeter Array. They determined, using continuum emission images, that MM1 remains unresolved with a size of $\\sim$0.03 pc and a mass of 29 \\Msun. Their molecular line spectrum shows the presence of several complex molecules, suggesting that MM1 is a hot core. They also found an extended \\tco\\ structure which may be evidence of a rotating envelope surrounding the central high--mass star, perpendicular to one of the outflows that Shepherd et al. (2007) detected toward MM1.  In this work we report observations of 11 molecular transitions using the APEX 12-m, SEST 15-m, and Nobeyama 45-m telescopes toward the G34.43+0.24 cloud. Since different molecular species trace different density and kinematics regimes, the capability of observing multiple spectral lines is particularly advantageous to determine the morphology and kinematics of star forming cores. Due to the high abundance of CO, the \\dco\\ line is an ideal tracer of the low intensity, high--velocity wing emission. The \\tco\\ line, being more optically thin than \\dco, traces, in addition to the ambient gas, the low--velocity gas of the molecular outflow. The \\sio\\ line is usually associated with the presence of outflows since its abundance is highly enhanced by shocks. Methanol transitions are usually used as tracers of dense gas, temperature probes and tracers of shocks, since this molecule is evaporated from grain mantles. The \\hco\\ and \\csdos\\ lines are high--density tracers and thus are good probes of the dense envelopes surrounding protostars. The \\cdo, \\cdodos, \\hcotrece, \\cssiete\\ and \\cstyc\\ molecular transitions probe the inner regions of the cores. The last two species are weakly influenced by infall and outflow motions, and are used as reference to estimate the central ambient cloud velocity.  The main goals of these multi--line observations are to determine the morphology and kinematics of the gas within the molecular cores in the G34.43+0.24 filamentary IRDC, in particular to investigate the presence of outflowing gas, in order to increase our knowledge of the evolutionary state of the cores and to achieve a better understanding of the overall process of massive star formation along the filament. This is the first survey of several molecular lines mapping portions and the whole cloud. It gives us, for the first time, a comprehensive view of the molecular gas distribution at the scale of parsecs. ",
        "conclusions": "We carried out a multi--line study of cores within the filamentary IRDC  G34.43+0.24. The emission from the filament was fully mapped in the \\dco\\ and \\tco\\ lines. In addition, emission from the central region of the filament was observed in eight different molecular transitions [\\cdo, \\cssiete, \\csdos, \\cstyc, \\hco, \\hcotrece, \\sio\\ and \\metanol], and from the north region in the \\csdos\\ and \\cdo\\ lines. The results of this study are summarized as follows.  The spatial distribution of the molecular emission in all, except methanol, lines is similar to that of the dust continuum emission derived from observations with similar angular resolution, showing a filamentary structure and four cores.  The mass of the molecular cores derived assuming LTE conditions are 300 and 1460 \\Msun\\ for the MM1 and MM2 cores, respectively. The masses derived assuming virial equilibrium are 1100, 1500 and 1400 \\Msun\\ for the MM1, MM2 and MM3 cores, respectively. The average molecular hydrogen densities of the cores are $1.8\\times 10^5$ cm$^{-3}$ in MM1, $1.7\\times 10^5$ cm$^{-3}$ in MM2 and $1.9\\times10^4$ cm$^{-3}$ in MM3. We also find that the molecular hydrogen column densities at the peak positions of MM1 and MM2 are, respectively, $2\\times 10^{22}$ cm$^{-2}$ and $8\\times 10^{22}$ cm$^{-2}$.  The derived parameters of the MM1 and MM2 cores are typical of massive and dense cores harboring high-mass YSO's (Plume et al. 1997; Garay et al. 2007). Those of the MM3 core suggest that it harbors young intermediate mass stars, in accord with the result obtained by Shepherd et al. (2007) from the modeling of the SED of the NIR sources detected using {\\it Spitzer}.  We modeled the molecular emission in the \\tco, \\cdo, and \\cssiete\\ lines toward the MM2 core, concluding that this core is undergoing infalling motions, with an infall velocity of $\\sim$1.3 \\kms\\ and a mass infall rate of $\\sim$$1.8\\times 10^{-3}$ \\Msun\\ yr$^{-1}$. This value is large enough to allow the formation of massive stars by accretion (see Osorio, Lizano \\& D'Alessio 1999).  We report the discovery of a molecular outflow associated with the massive, dense MM3 core located in the northern region of G34.43+0.24, and possibly of a second one associated with the MM4 core. The molecular outflow within MM3 has a total mass of 41 \\Msun, a momentum of 270 \\Msun\\ \\kms, and a kinetic energy of 1120 \\Msun\\ km$^2$ s$^{-2}$; indicating that it is driven by high--intermadiate mass stars. We also detected high velocity molecular gas toward the MM1 and MM2 cores with total masses of 25 and 80 \\Msun, momenta of 230 and 690 \\Msun\\ \\kms, and kinetic energies of 1290 and 3630 \\Msun\\ km$^2$ s$^{-2}$, respectively."
    },
    "1005/1005.3562_arXiv.txt": {
        "abstract": "We performed stellar population synthesis on the nuclear and extended regions of NGC~1068 by means of near-infrared spectroscopy to disentangle their spectral energy distribution components. This is the first time that such a technique is applied to the whole 0.8$-$2.4~$\\mu$m wavelength interval in this galaxy. NGC~1068 is one of the nearest and probably the most studied Seyfert 2 galaxy, becoming an excellent laboratory to study the interaction between black holes, the jets that they can produce and the medium in which they propagate. Our main result is that traces of young stellar population are found at $\\sim$100~pc south of the nucleus. The contribution of a power-law continuum in the centre is about 25$\\%$, which is expected if the light is scattered from a Seyfert 1 nucleus. We find peaks in the contribution of the featureless continuum about 100 - 150 pc from the nucleus on both sides. They might be associated with regions where the jet encounters dense clouds. Further support to this scenario is given by the peaks of hot dust distribution found around these same regions and the H$_2$ emission line profile,  leading us to propose that the peaks might be associate to regions where stars are being formed. Hot dust also has an important contribution to the nuclear region, reinforcing the idea of the presence of a dense, circumnuclear torus in this galaxy. Cold dust appears mostly in the south direction, which supports the view that the southwest emission is behind the plane of the galaxy and is extinguished very likely by dust in the plane. Intermediate age stellar population contributes significantly to the continuum, specially in the inner 200~pc. ",
        "introduction": "The coexistence of black holes and starburst clusters is known to exist in many galaxies, and there are many evidences that suggest a connection between these phenomena. Many studies point out that both the active nucleus and starbursts might be related to gas inflow, probably triggered by an axis-asymmetry perturbation like bars, mergers or tidal interactions (Shlosman, Frank \\& Begelman 1989, Shlosman, Begelman \\& Frank 1990, Maiolino et al. 1997, Knapen, Shlosman \\& Peletier 2000, Fathi et al. 2006, Riffel et al. 2008). In addition, one of the most intriguing research areas in contemporary extragalactic astrophysics involves the study of the interplay between nuclear black holes, the jets which they can produce and the interstellar/intergalactic medium (ISM) in which they propagate. These jets can have a considerable impact on this medium. One aspect of jet-ISM interaction is that it can trigger star formation. Such jet-induced star formation is considered a possible mechanism to explain the UV continuum emission observed in the host galaxies of distant radio sources and the ``alignment effect'' between the radio emission and this continuum (Rees 1989). Although this effect might play a very important role in high-z radio galaxies, detecting and studying the jet-ISM interaction in them is very challenging. Because of the observational problems, it is important to find nearby examples of this kind of interaction where a detailed study can be carried out.  NGC 1068 is an ideal object in this case. It is one of the nearest and probably the most intensely studied Seyfert 2 galaxy. Observations in all wavelength bands from radio to hard X-rays have formed a uniquely detailed picture of this source. NGC 1068 hosts a prominent narrow-line region (NLR) that is approximately co-spatial with a linear radio source with two lobes (Wilson \\& Ulvestad 1983). Star formation activity coexistent with the active galactic nucleus (AGN) was detected on both larger (e.g., Telesco \\& Decher 1988) and smaller scales (Macchetto et al. 1994; Thatte et al. 1997). However, the link between all the processes is still under debate.  The Near Infrared Region (NIR) is particularly interesting to help unveiling this link because it is accessible to ground-based telescopes and, at the same time, able to probe highly obscured sources. However, tracking the star formation in the NIR is not simple (Origlia \\& Oliva 2000) although recent studies exploring this region have already shown its strong potential at detecting intermediate-age stellar population not easily tracked in the optical without ambiguity (e.g. Riffel et al. 2009, Davies et al. 2007). At near-IR wavelengths stellar photospheres usually remain the dominant sources of light, and galaxy spectra are shaped by red supergiants (RSG) shortly after starbursts, and then by giants of the first and of the asymptotic giant branches (AGB).  AGB stars are rare members of stellar populations. However, they are among the most luminous cool stars and can therefore be detected sometimes even individually in galaxies. The TP-AGB stars leave a unique fingerprint on the integrated spectra, like the 1.1 $\\micron$ CN band (Maraston 2005, Riffel et al. 2007, 2009). Hence when detected, they can help to determine the age of the stellar population through the integrated light. The contribution of this stellar phase in stellar population models has been recently included in both the energetics and the spectral features (Maraston 2005). In particular these models employ empirical spectra of oxygen-rich and carbon stars (Lan\\c con \\& Wood 2000), which are able to foresee characteristic NIR absorption features.  With this in mind, we present here for the first time in the literature a detailed fitting of the continuum emission components in the 0.8$-$2.4 $\\micron$ interval of NGC~1068 across the central 15\" ($\\sim$ 1100 pc) of this source. The main purpose is to determine the fraction with which the different components contributes to the observed integrated light and how they are related to each other. The paper is structured as follows: in \\S 2 we describe the observations. In \\S 3 we describe the fitting method and in \\S 4 the results are presented and discussed. Final remarks are given in \\S 6. ",
        "conclusions": "We investigated the NIR spectra of NGC~1068 extended across 15\" ($\\sim$ 1100 pc) in the North-South direction. The spectra was taken with the IRTF SpeX instrument, obtained in the short cross-dispersed mode. The objective was to disentangle the galaxy's spectral energy distribution components along the full wavelength coverage (0.8 $\\micron$ - 2.4 $\\micron$) using stellar population synthesis. This was done here for the first time for this galaxy. We used the STARLIGHT code, which considers the whole observed spectrum, continuum and absorption features.  We find that intermediate age stellar population contributes significantly to the continuum, specially in the inner 200 pc. We don't have the spatial resolution to determine the size of this nuclear stellar cluster, but results agree with previous determinations of the stellar population of NGC~1068 from Davies et al. (2007) and Riffel et al. (2009).  We find a small contribution of a young stellar population component about 100 pc from the nucleus. Although this result has to be considered carefully, since the detection might be an artefact from the method used, taken together with our other findings this suggests that this is the region where the jet interacts with the ISM and might be forming stars.  The contribution of a power-law continuum in the centre is about 25$\\%$, which is expected from a scattered Seyfert~1 nucleus (Cid Fernandes et al. 1995). We also find peaks in the contribution of the featureless continuum about 100 - 150 pc from the nucleus, both in the north and in the south direction. These might be associated with the region where the jet encounters dense clouds. This result is further supported by peaks of hot dust found around the same regions. We compared these results with the H$_2$ emission line profile obtained from Paper I. This line also shows extended emission around the same region where the hot dust is found. Taking this together with the young stellar population detection, we believe there is strong evidence that the interaction of the jet with the dense ISM in NGC~1068 is forming stars. Further investigation with more spatial resolution is needed to confirm this hypothesis.  Hot dust also has an important contribution to the nuclear region, which reinforces the idea of the presence of a dense, circumnuclear torus in this galaxy. Cold dust appears mostly in the south direction, which supports the idea that the southwest emission is behind the plane of the galaxy, and is extinguished by the dust in the plane. We determine the dust mass in this galaxy to be about 3.4 $\\times$ 10$^{-2}$ M$_{\\sun}$, which puts NGC 1068 as one of the Seyferts with the highest dust content measured."
    },
    "1005/1005.5422_arXiv.txt": {
        "abstract": "Much of the progress in our understanding of dynamo mechanisms has been made within the theoretical framework of magnetohydrodynamics (MHD). However, for sufficiently diffuse media, the Hall effect eventually becomes non-negligible. We present results from three dimensional simulations of the Hall-MHD equations subjected to random non-helical forcing. We study the role of the Hall effect in the dynamo efficiency for different values of the Hall parameter, using a pseudospectral code to achieve exponentially fast convergence. We also study energy transfer rates among spatial scales to determine the relative importance of the various nonlinear effects in the dynamo process and in the energy cascade. The Hall effect produces a reduction of the direct energy cascade at scales larger than the Hall scale, and therefore leads to smaller energy dissipation rates. Finally, we present results stemming from simulations at large magnetic Prandtl numbers, which is the relevant regime in hot and diffuse media such a the interstellar medium. ",
        "introduction": "\\label{sec:intro} The generation of magnetic fields by dynamo activity plays an important role in a wide range of astrophysical objects, ranging from stars to clusters of galaxies. The gas in these objects is characterized by turbulent flows, as shown for instance by scintillation observations of the interstellar medium \\cite{Span01, Mint96}, or from pressure maps in galaxy clusters \\cite{Schue04}. Mechanisms able to generate magnetic fields by dynamo action are often classified as large- and small-scale dynamos, depending on the correlation length of the induced magnetic field. In this context, large and small are referred to the energy containing scale of the turbulent hydrodynamic flow. This classification is not rigid, as in many astrophysical objects both dynamos may be at work, but it gives a useful framework considering the limitations in the scale separation that can be achieved in numerical simulations. Also, the physical properties of the flows that can give rise to one or the other are somewhat different.  Helical flows have proved efficient in generating large-scale dynamos, i.e., on scales larger than the energy-containing eddies of the flow \\cite{Pouq76,Meneg81,Brand01,Gomez05}. It is now known that large-scale dynamo action can also be produced by anisotropic and inhomogeneous flows (e.g., flows with a large scale shear).  On the other hand, non-helical flows can be instrumental in generating small-scale dynamos \\cite{Kazantsev68}, i.e., on sizes smaller than those of the energy-containing eddies \\cite{Scheko01,Scheko04a,Haugen04}. In recent years, the study of small-scale dynamos with magnetic Prandtl number $Pm = \\nu /\\eta $ (the ratio between the viscosity and the magnetic diffusivity of the plasma) different from unity has received special attention \\cite{Scheko04b,Mininnietal05}, both for $Pm \\gg 1$ \\cite{Scheko02} and for $Pm \\ll 1$ \\cite{Ponty05,Iska07}.  Motivations to study these regimes include recent experiments of dynamo action using liquid sodium \\cite{Monchaux07}, as well as the fact that many astrophysical plasmas are characterized by magnetic Prandtl numbers different from unity. For instance, the magnetic Prandtl number is much smaller than one in the solar convective region, and it is typically much larger than one in the interplanetary medium and also in the interstellar medium (ISM).  For sufficiently low-density media such as the one that pervades the ISM, kinetic effects such as the Hall effect or ambipolar diffusion might also become relevant \\cite{Sano02}. The potential relevance of ambipolar diffusion in astrophysical dynamos was studied in Refs.~\\cite{Branden00,Zweibel02}. The relevance of the Hall effect has been recognized in various astrophysical applications \\cite{Balbus01,Sano02,Mininni02}, space plasmas \\cite{Deng01,Oieroset01,Mozer02}, and also laboratory plasmas \\cite{Yamada97,Mirnov03,Ji08}. The role of the Hall effect on large-scale dynamos subjected to helical forcing has also been addressed in the literature \\cite{Mininni03a,Mininni03b}. Less attention has received the impact of kinetic effects on the small-scale dynamo. A theoretical model of the kinematic small-scale dynamo with Hall effect was presented in \\cite{Kleeorin94}, but to the best of our knowledge no numerical studies of the non-linear and saturated regime were considered in the literature.  In this paper, we present results from three dimensional simulations of the Hall-MHD equations subjected to random non-helical forcing. The main aim is to study the role of the Hall effect in the small-scale dynamo efficiency for different values of the Hall parameter. As a result of the study, we also discuss the impact of the Hall effect on the dynamo saturation values, and on magnetic and total dissipation rates. The structure of the paper is as follows.  A brief introduction to the theoretical framework known as Hall-MHD is presented in Sect.~\\ref{sec:hmhd}. The role of the Hall effect in the efficiency of the dynamo is shown in Sect.~\\ref{sec:lin}. In Sect.~\\ref{sec:spec} we characterize the stationary regime that is attained when the dynamo process saturates, showing the corresponding energy power spectra. The energy transfer rates participating in the nonlinear energy cascade are displayed in Sect.~\\ref{sec:trans}. In Sect.~\\ref{sec:Pm}, we explore the regime of large magnetic Prandtl number (i.e., when the viscous dissipation scale is larger than the resistive dissipation scale) which, as mentioned, is particularly relevant in diffuse media such as the ISM. Finally, the conclusions are summarized in Sect.~\\ref{sec:conclu}. ",
        "conclusions": "\\label{sec:conclu}  We present results from three dimensional simulations of small-scale dynamo action for magnetic Prandtl numbers $Pm=1$ and $10$ in conducting flows with the Hall effect. This effect is believed to be non-negligible in sufficiently diffuse media, and its relevance has been recognized in various astrophysical, space, and laboratory plasmas. As a first step toward a better description of dynamo action in such media, only the incompressible Hall-MHD equations were solved, and the inclusion of compressible effects as well as other kinetic effects such as ambipolar difussion is left for future studies.  However, the inclusion of only the Hall effect acting at the smallest relevant dynamical scales of the flow gives rise to measurable differences with previous studies of dynamo action. A magnetic non-linear regime is identified when the magnetic field (and the current density) becomes large enough to differentiate the electron velocity from the bulk flow velocity. After saturation, differences in the stationary level of magnetic energy and in the total and magnetic energy dissipation rates are obtained, depending on the amplitude of the Hall effect. Finally, the peak of the current density spectrum is found to be dependent on the strength of the Hall term, with its peak moving toward larger scales (smaller wavenumbers) as the Hall scale is increased.  By studying the detailed transfer of energy among fields and scales, we observe that the effect of the Hall term is twofold: it transfers energy towards larger scales for scales larger than the Hall length, and it transfers energy towards smaller scales for scales smaller than this length. The modification of the energy flux resulting from this transfer is consistent with the observed changes in the saturation values of energy and dissipation rate observed in our simulations."
    },
    "1005/1005.2202_arXiv.txt": {
        "abstract": "{The \\object{NGC\\,1999} reflection nebula features a dark patch with a size of $\\sim$10,000\\,AU, which has been interpreted as a small, dense foreground globule and possible site of imminent star formation.  We present \\herschel{} PACS far-infrared 70 and 160\\,\\micron{} maps, which reveal a flux deficit at the location of the globule.  We estimate the globule mass needed to produce such an absorption feature to be a few tenths to a few \\msun. Inspired by this \\herschel{} observation, we obtained APEX LABOCA and SABOCA submillimeter continuum maps, and Magellan PANIC near--infrared images of the region.  We do not detect a submillimer source at the location of the \\herschel\\ flux decrement; furthermore our observations place an upper limit on the mass of the globule of $\\sim$2.4$\\cdot 10^{-2}$\\,\\msun. Indeed, the submillimeter maps appear to show a flux depression as well. Furthermore, the near--infrared images detect faint background stars that are less affected by extinction inside the dark patch than in its surroundings. We suggest that the dark patch is in fact a hole or cavity in the material producing the \\ngc{} reflection nebula, excavated by protostellar jets from the V\\,380\\,Ori multiple system. } ",
        "introduction": "In the year 1774 C.E.\\ Sir Friedrich Wilhelm Herschel first noticed patches of sky in the constellation Scorpio that were devoid of stars. Unable to find even the faintest star in these regions, his sister Caroline reported him to exclaim: ``Hier ist wahrhaftig ein Loch im Himmel!'' (``Truly there is a hole in the sky here!'').  These dark areas are now known to be due to obscuring material \\citep[e.g.,][]{barnard27}: clouds of molecular gas, whose dust content absorbs the light of background stars.  Dark clouds, ranging in mass and size from giant molecular clouds (GMCs) to tiny globules, are the sites of star formation in our Galaxy.  \\ngc{} is a small reflection nebula in the \\object{LDN\\,1641} portion of the Orion\\,A GMC, located in a small group of 22 pre--main sequence stars and protostars, including the driving source of the prototypical Herbig--Haro objects \\object{HH\\,1} and \\object{HH\\,2} (Megeath et al., in prep.).  It is illuminated by the Herbig\\,Ae/Be star \\object{V\\,380\\,Ori}, a multiple system with a circumsystem disk where the primary is a 100\\,L$_{\\odot}$ B9 star exhibiting strong emission lines \\citep[e.g.,][]{leinertetal1997,hillenbrandetal1992,alecianetal2009}. \\ngc{} features a compact (20--30\\arcsec{}) dark patch (see Fig.~\\ref{fig:ngc1999_hst}), which is described in the literature as a dark globule and potential site of star formation \\citep[e.g.,][]{herbig1946,herbig1960,warrensmithetal1980}, but which has never been studied in greater detail before.  \\begin{figure} \\begin{center} \\scalebox{0.325}{{\\includegraphics[angle=270]{14612fg1.ps}}} \\caption{\\ngc{} and HH\\,1/2 region (DSS, left) and HST F450W/F555W/F675W true color image of the \\ngc{} dark patch. \\label{fig:ngc1999_hst}} \\end{center} \\end{figure}   We report here the serendipitous and surprising detection with PACS of a 70 and 160\\,\\micron{} flux decrement at the location of \\ngc{}. The maps show a dark patch against the nebulous far--IR emission that closely resembles the dark patch seen at visible wavelengths.  The MSX, ISO, and Spitzer space telescopes have detected clouds in absorption at mid--IR (5-30\\,\\micron{}) wavelengths against the diffuse IR background of the Galaxy \\citep[e.g.,][]{bacmann00,stutz08,tobin10,stutz09a}, in a few cases even at 70\\,\\micron{} \\citep{stutz09b}.  However, the detection of the \\ngc{} globule by PACS would be the first detection of a dark globule in absorption at 160\\,\\micron{}. At wavelengths $\\ge 160$\\,\\micron{}, cold globules are normally detected through the emission from cold (10--30\\,K) dust.  Furthermore, absorption at these wavelengths requires extinctions in excess of 100\\,$A_V$.  Motivated by the possible discovery of a 160\\,\\micron{} dark cloud, we obtained follow--up ground based submillimeter and near--IR (extinction) observations towards the globule.  We did not detect the column density or mass of cold dust necessary to produce such an absorption feature in the far--IR.  We therefore conclude that the dark globule is actually a cavity in the \\ngc{} cloud carved by outflows from the V\\,380\\,Ori system.  Thus, the \\herschel{} telescope has discovered what truly is a hole in the sky. ",
        "conclusions": "Optical images of \\ngc{} show a dark patch suggestive of a small, dense globule obscuring the \\ngc{} reflection nebula. Surprisingly, our 70 and 160\\,\\micron{} images clearly show a dark patch against the bright nebula with a strikingly similar morphology to that in visible light images.  These \\herschel{} measurements, combined with subsequent ground based data, lead to the conclusion that the dark feature is a hole in the \\ngc{} nebula. First, we find that the masses needed to cause the 70 and 160\\,\\micron{} dark patch --- 0.1 and 2.5\\,\\msun{}, respectively --- are inconsistent. Furthermore, the globule is not detected in emission at 350 and 870\\,\\micron{}: the SABOCA obervations place an upper limit of $2.4\\cdot 10^{-2}$\\,\\msun{} for a temperature of 10\\,K.  This upper limit is far below the amount of mass needed to cause the obscuration in the PACS data.  Finally, near--IR observations with PANIC detect background stars toward the  dark patch; the $H-K_\\mathrm{s}$ colors of these stars are slightly bluer than stars detected outside the globule, suggesting a lower extinction toward the dark patch. Furthermore, the extinctions of the background stars are less than that required to absorb the 160\\,\\micron{} flux. Taken together, these observation show that the dark patch is not a globule, but instead a cavity in the nebula.  The presence of a well delineated cavity of the size of $\\sim$10,000\\,AU deserves some attention. With the typical turbulent velocities on the order of a few km/s in clouds, such a cavity should be filled on timescales of at most a few 10,000\\,yrs and quickly disappear. The PANIC H$_2$ narrow band data (Fig.\\,\\ref{fig:NGC1999_H2}) deliver important hints about the possible origin and peculiar shape of the cavity. They reveal a previously unknown, faint H$_2$ bow shock enveloping the SMZ\\,60/HH\\,148 compact knots \\citep{stankeetal2002,corcoranray1995}, constituting the clearest evidence so far for a collimated flow running northeast to southwest through the cavity. This flow, which likely originates in the V\\,380\\,Ori multiple system, could possibly excavate the southern part of the cavity. \\object{SMZ\\,6-8} in the southeastern corner of the PANIC image resembles a small bow shock in a flow coming from the northwest; together with \\object{HH\\,35}, located northwest of V\\,380\\,Ori (Fig.\\,\\ref{fig:ngc1999_hst}), it indicates a second, northwest to southeast oriented flow, which could be responsible for digging the northwestern lobe of the cavity.  We note a similar flux depression in the 160\\,\\micron{} \\herschel{} image southwest of the protostar HOPS\\,166, which drives the \\object{HH\\,147} outflow \\citep{corcoranray1995}. Optical images (e.g., Fig.\\,\\ref{fig:ngc1999_hst}) show a circular reflection nebula marking the HH\\,147 outflow cavity, coinciding with the far--IR flux depression. The \\ngc{} dark patch may therefore only be a somewhat peculiar example of a cavity carved in the ambient medium by an outflow, rendered particularly visible by the illumination and heating of the cavity walls by V\\,380\\,Ori. Sensitive far--IR maps taken with \\herschel{} may therefore provide a new tool to assess the importance of outflow feedback on cloud cores."
    },
    "1005/1005.2943_arXiv.txt": {
        "abstract": "In a series of papers, we propose a theory to explain the formation and properties of rings and spirals in barred galaxies. The building blocks of these structures are orbits guided by the manifolds emanating from the unstable Lagrangian points located near the ends of the bar. In this paper, the last of the series, we present more comparisons of our theoretical results to observations and also give new predictions for further comparisons. Our theory provides the right building blocks for the rectangular-like bar outline and for ansae. We consider how our results can be used to give estimates for the pattern speed values, as well as their effect on abundance gradients in barred galaxies. We present the kinematics along the manifold loci, to allow comparisons with the observed kinematics along the ring and spiral loci. We consider gaseous arms and their relations to stellar ones. We discuss several theoretical aspects and stress that the orbits that constitute the building blocks of the spirals and rings are chaotic. They are, nevertheless, spatially well confined by the manifolds and are thus able to outline the relevant structures. Such chaos can be termed `confined chaos' and can play a very important role in understanding the formation and evolution of galaxy structures and in galactic dynamics in general. This work, in agreement with several others, argues convincingly that galactic dynamic studies should not be limited to the study of regular motions and orbits. ",
        "introduction": "\\label{sec:intro}  In previous papers (\\citealt{RomeroGMAG06}, hereafter Paper I; \\citealt{RomeroGAMG07}, hereafter Paper II; \\citealt{AthaRGM09}, hereafter paper III; \\citealt{RomeroGMGA09}) we proposed a theory to explain the formation and structure of spirals and inner and outer rings in barred galaxy potentials. We propose that the building blocks of these structures are chaotic orbits that are guided by manifolds associated with the Lagrangian points $L_1$ and $L_2$ which are located along the direction of the bar major axis. These manifolds can be thought of as tubes that confine the orbits, so that the latter can form thin structures in configuration space.  A theory, however, can be dynamically correct but still irrelevant to a particular application. We therefore have to check whether our theory is applicable to spirals and rings observed in barred disc galaxies. It is thus necessary to compare our theoretical results and predictions to observations. We started this in Paper IV (\\citealt{AthaRGBM09}, hereafter Paper IV) where we made a number of comparisons of model spirals and rings, the latter both inner and outer, to observations and found good agreement. We found that our theory can reproduce all necessary morphologies of both inner and outer rings, and produces no unrealistic morphologies. Model inner rings were found to be elongated along the bar and outer ones either along it ($R_2$, or $R_2'$), or perpendicular to it ($R_1$, or $R_1'$), or both ($R_1R_2$).  Model spirals in barred galaxy potentials were found to be predominantly two-armed and trailing, although arms of higher multiplicity are possible for specific potentials. They were found to have the right shape and could reproduce the fall-back of an arm towards the bar region or towards the other arm, which is observed in many spirals. A quantitative comparison to the arm shapes of NGC 1365 proved successful. We predict that the relative strength of the non-axisymmetric forcing in the region around and somewhat beyond corotation influences the winding of the arms, in the sense that in strongly barred galaxies the spirals will be more open than in less strongly barred ones. Thus, a series of models with increasing bar strength will have a continuous sequence of morphologies from $R_1$ to $R_1'$, then to tightly wound spirals, to end with open spirals. We also compared the shape of the inner and outer rings, as well as the ratio of outer-to-inner ring major axes of our models to observations and discussed which type of potentials give best agreement. We also showed that there are correlations, or trends, between all the ratios of ring sizes discussed above and the bar strength. These correlations are very tight if only one type of model potential is used but thicken, sometimes very considerably, if other models, with different properties and analytical expressions, are included.  The present paper is the fifth and last of this series. A reminder of the main theoretical results and prerequisites is given in Sect.~\\ref{sec:theory}. This is very brief, since these concepts have been described extensively in Papers I, II and III. Here we will mainly present further comparisons with observations and predictions. These include e.g. the bar shape (Sect.~\\ref{subsec:barshape}), ansae (Sect.~\\ref{subsec:ansae}), radial drift and abundance gradients (Sect.~\\ref{subsec:abundancegr}). Kinematics and behaviour of the line-of-sight velocities are discussed in Sect.~\\ref{sec:Kinematics} and pattern speed prediction is the subject of Sect.~\\ref{sec:omegap}. In Sect.~\\ref{sec:discussion} we discuss specific topics, such as time evolution, self-consistency and spiral structure formation, and we also compare with other spiral structure theories and give outlines for further work. These discussions rely on material from all five papers in this series. Finally, we summarise and conclude in Sect.~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  Since the early work of B. Lindblad (\\citeyear{Lindblad63}), density waves have been commonly assumed to be at the basis of any explanation of spiral structure formation. In this series of papers we presented an alternative viewpoint, which we have applied specifically to barred galaxies. This explains the formation of spiral arms, as well as of inner and outer rings, in a common theoretical framework. According to our theory, it is the unstable Lagrangian points located at the ends of the bar and the corresponding manifolds that are responsible for the formation of spirals and rings. These manifolds drive orbits, which are in fact chaotic, but are confined by the manifolds, so that they create over-densities which have the right shape to explain the spirals and the rings (Papers I and II).  We showed that weaker non-axisymmetric perturbations will produce manifolds of $R_1$ ring shape, while stronger ones will produce other types of rings and spirals. It is, nevertheless, the same theory that can explain all these different morphologies. This is in good agreement with the observed morphological continuity between rings, pseudorings and spirals. This continuity is clear in our theory, but not in others.  We made many different kinds of comparisons with observations. We first made sure that our theory can account for the basic morphological structures observed, namely the spirals, the inner rings and the outer rings (Paper IV). The model outer rings have the observed $R_1$, $R_2$ and $R_1R_2$ ring and pseudo-ring size and morphologies, including the dimples often observed near or at the direction of the bar major axis. Both model and observed inner rings have roughly the same relative sizes and orientations.  We next turned to the main spiral arm properties (Paper IV). We showed that, as the bar grows, material gets trapped mainly in the unstable manifolds, i.e. the sense of the arms will be trailing, in good agreement with what is currently known from observations. There is, nevertheless, a very small fraction of the material concerned that will be caught by the stable manifold branch, and will thus form a very low amplitude leading spiral. If its amplitude is not too low, it may be observable in a Fourier analysis of the galaxy image, or as interference patterns on the arm amplitude.  Theoretical and observational results agree also well on the number of arms in a barred galaxy (Paper IV). Our theory links one spiral arm to each of the unstable Lagrangian points in the standard case. Since galactic bars are bisymmetric, there should be two such Lagrangian points, the $L_1$ and $L_2$, and therefore two spiral arms. This morphology persists even in the non-standard case where the $L_1$ and $L_2$ are stable (Paper III). Indeed, the vast majority of barred galaxies have two spiral arms. There are, nevertheless, a few barred galaxies with four spiral arms. Our theory can account for such cases, given the right potential, and we discussed a few such potentials in Paper IV.  We also compared the shape of observed spiral arms with those of manifolds (Paper IV). We showed that our theory can explain the characteristic arm winding often observed, namely that, as the angle along the arm increases, the radius first increases and then, after reaching a maximum, `falls back' towards the bar. This property has, to our knowledge, not been reproduced by any other theory so far, but follows naturally from ours. We also predict that spirals in galaxies with stronger bars will be more open than those in galaxies with weaker bars. This prediction can be made because our theory, contrary to most others, is non-linear. Finally, very tightly wound near-logarithmic spirals can also  be obtained with our theory, but will rely considerably on spiral forcing.  The shape of rings and the ratio of their extents, although not as straightforward to compare as other properties, reveal interesting information (Paper IV). Concerning $R_1$ rings, for which we have sufficient models to get adequate statistics, we find that the ring shapes cover the same range as observations. A comparison for inner rings is less straightforward, because in strong bar cases parts of the inner manifold branches outline parts of the bar and not an inner ring. Taking this into account, we find that there is a fair agreement between models and observations.  One of the predictions made in Paper IV concerns the shape of the inner rings. Namely, we found a strong correlation between the ring shape and $Q_{t,L_1}$, in the sense that bars with a stronger forcing at or somewhat beyond $L_1$ should have more elongated $R_1$ rings. We also stressed that, in order for correlations concerning a spiral or a ring property with the strength of the non-axisymmetric forcing to be strong, the latter quantity should be measured at or beyond the Lagrangian radius, i.e. with $Q_{tL_1}$. If $Q_b$ is used instead, the correlation can be destroyed, or severely damaged. Both of these results have now been confirmed by observations \\citep{GrouchyBSL10}.  We also introduced collisions and dissipation to the manifold calculations, in order to roughly model the gas properties (Paper IV). We found that this does not influence the existence of the spiral arms or rings, and makes no major modifications to their shape. The amount of dissipation, however, does influence the width of the arms. These become thinner as the dissipation is increased, so that the gaseous arm comes nearer to the lowest energy manifold.  Our theory makes very clear and, for many morphological types, precise predictions about the value of the bar pattern speed, since it shows how it is possible from the morphology only to locate the positions of the $L_1$ and $L_2$, both for spirals and for the various ring morphologies (Paper V). This of course works only in the case where sufficient morphological features are present in the observed galaxy. For example if there is only a bar, with no spirals or rings, our theory can not be applied and the $L_1$ and $L_2$ can not be located. However, in cases with a reasonable amount of structure our theory allows the location of the $L_1$ and $L_2$ from morphology alone. In cases with appropriate morphological features this may allow a more accurate determination of the pattern speed then other methods used so far. Our theory also provides building blocks to explain the rectangular-like outline observed in strong barred galaxies and to explain ansae (Paper V).  Concerning photometry, our results are not very constraining. The amplitude of the spiral should in general decrease outwards, but this decrease may well not be monotonic, because other arm components -- e.g. a weak, leading component --, if present, will lead to bumps on an otherwise smoothly decreasing density profile.  The circulation of material within the manifolds also depends on the bar strength. In the relatively weak bars, which form $rR_1$ morphologies, the mass elements can move from the inner manifolds to the outer manifolds and then back to inner ones again. I.e. material can move from the region within corotation to the region outside it and vice versa via the neighbourhood of the $L_1$ and $L_2$. Averaged over time, this brings internal circulation, but no net motion of material inwards or outwards. The situation is more complex for the cases where the non-axisymmetric forcing around corotation is stronger and which have a spiral morphology. In such cases, material moves from the region within corotation to the region outside it. If the potential has a predominantly barred structure, this material returns eventually to the region of the bar or of the other arm. But if the potential has also a spiral component, then the arm can continue winding outwards for much longer times. Either way, this results in a considerable radial mixing and material can reach the outermost parts of the disc and can contribute to its radial extension. Of course the density wave theory had already predicted that in a spiral galaxy the radial motion is outwards outside corotation \\citep{Lin.Shu71, Kalnajs78}. Our results, however, are more than a simple restatement of the above, because our theory shows how material from well within corotation can reach regions well outside it. It can thus explain considerable radial mixing and have a more important impact on the formation of the outer disc than those of the density wave theory. Furthermore, our theory predicts the relation between bar strength and the abundance gradients found in observations \\citep{Martin.Roy94}.  We also made predictions on the galaxy kinematics, by measuring the radial and tangential velocity components along the manifold loci, i.e. along the spirals and the inner and outer rings. For the tangential component we subtract the local circular velocity and thus obtain the perturbation, or peculiar tangential velocity. Plotted as a function of the azimuthal angle along the ring, both radial and tangential velocities show $m$ = 2 oscillations, the amplitude of the former being higher than that of the latter. We find that the amplitude of these oscillations increases with increasing bar strength and that they are larger in the inner than in the outer ring. The latter result is in agreement with the former, since the bar forcing is higher at the inner ring than at the outer one.  For spirals, we find that the radial velocity of material moving from the Lagrangian points outwards along the arm increases with the distance to the Lagrangian point until it reaches a maximum and then it decreases, until at some point it changes sign. From that point onwards the arm will start approaching the other arm or the region of the bar. Where this occurs is a function of the bar strength and, when observed, will provide an additional confirmation to our theory. Nevertheless, this may be a difficult observation to make, since the spiral density is low in that region. Furthermore, in galaxies with a high spiral amplitude, the spiral forcing may be a considerable fraction of the non-axisymmetric forcing, so that this sign reversal may not take place. Finally, we calculated the line-of-sight velocity and found that in the case of rings its variation along the ring can be well modelled by a sinusoidal. This is not true for very strong bars and spiral morphologies. Finally, when the sinusoidal is a reasonable fit, we find that the difference between the kinematic and photometric major axes is in agreement with the observed values.  As already mentioned, our models show a clear connection between bar strength and morphology of the response. Relatively weak non-axisymmetric forcings give $R_1$ and $R_1'$ morphologies, while stronger forcings give spirals and other types of rings. Yet it is possible in self-consistent calculations, in which the potential due to the material trapped in the manifold is taken into account, to form a spiral in cases where the bar on its own is too weak to sustain it. This is due to the fact that initially the only part of the ring that will be populated is the part which is near the $L_1$ or the $L_2$. This then adds a spiral-like forcing to that of the bar, so that manifold will have a more spiral-like shape. It is thus possible in self-consistent calculations to have a spiral morphology in a bar potential which is somewhat below the threshold between the $R_1$ and the spiral morphology in the rigid bar forcing case.  In these series of papers, we applied our theory to barred galaxies. It can, however, be straightforwardly applied also to non-barred galaxies with other non-axisymmetric perturbations, such as spirals or ovals, provided these have unstable saddle Lagrangian points. There is thus no discontinuity between barred and non-barred spiral galaxies. On the other hand it is unclear at this point whether, and under what conditions, complex unstable Lagrangian points can have the type of manifolds that can create realistic spirals or rings. Such Lagrangian points are found e.g. at the $L_4$ and $L_5$ Lagrangian points of very strong bars \\citep[][and Paper IV]{Atha92a}. More work is necessary before this question can be answered.  We discussed how material can get trapped in the manifolds, where the mass in the spirals comes from and whether the spirals are short-lived or long-lived. We also compared our theory with other theories not relying on chaotic orbits and manifolds. In particular, we discussed a number of advantages which our theory has over the density wave theory. We stressed that it need not necessarily be the same theory that explains all spirals in disc galaxies, and even in one single galaxy more than one mechanism can be at play. The properties of the resulting spiral would then result from the self-consistent interaction of all the spirals of different origin. We finally propose a number of avenues for future work within the framework of our theory.  We showed that the orbits that form the building blocks of the spirals and of the inner and outer rings are chaotic. Nevertheless, they can form narrow features because they are confined by the manifolds, which act as guiding tubes. We propose to call this type of chaos `confined chaos' and suggest that such orbits should be taken into account in dynamical studies, since they could contribute, together with the regular orbits trapped around the stable periodic orbits, to galactic structures.  A prerequisite for any theory is that it can be falsified \\citep[e.g.][]{Popper59}. Ours fulfils well this criterion. Indeed there are a number of observational results that could have proved it wrong. Namely, if the two armed global spirals were not preponderant in symmetric barred galaxies, or if these arms were leading, or if the major features of the observed and the theoretical morphologies were inconsistent, our theory would have been proven wrong. As we showed in this paper, this is not the case and there is good agreement between our theory and observations in these and many other points. There are more observations which could have proven our theory to be wrong. For example, our theory would be invalidated if it was clearly shown that all spirals rotate with a different pattern speed than the bar.  In total over the five papers, we presented a number of predictions of our theory. Whenever the necessary data are available, we compared the results of our theory to the existing observational results. All tests we have been able to make tend to show that our theory is viable. In many cases, however, new observations, or new analysis of existing data are necessary. In such cases, we presented here simply the theoretical predictions and outlined the way the comparisons could be made. The results of these comparisons will be very important in order to confirm our theory, to bring about modifications, or to reject it. As already mentioned, one of the theoretical predictions of paper IV concerning the dependence of the inner ring shape on bar strength has already been confirmed by observations \\citep{GrouchyBSL10}.  We can thus conclude that the theory we presented in this series of papers should be useful in order to explain the formation of spirals and of inner and outer rings in barred galaxies. Concerning rings, no other theory has been fully developed, while our results agree very well with those of simulations, suggesting that the motion of particles in them are guided by the manifolds we describe here. Concerning spirals, the comparisons we have made with observations are more extensive than those that have been made so far for the swing amplifier theory. Furthermore, our theory is clearly falsifiable, contrary to the classical WKBJ density wave theory (see discussion by \\cite{Kalnajs78}). We thus believe that it should be considered as a possible alternative, on a par with other theories. More work, both theoretical and observational, is necessary in order to establish in what cases each one of the so far proposed theories prevails and how they can, perhaps all together, come to  an explanation of spiral formation and spiral properties."
    },
    "1005/1005.2458_arXiv.txt": {
        "abstract": "The conditions of thermodynamic equilibrium, local thermodynamic equilibrium, and statistical equilibrium are discussed in detail. The equations of statistical equilibrium and the supplementary equations are shown together with the expressions for radiative and collisional rates with the emphasize on the solution for trace elements. ",
        "introduction": "Solution of the stellar atmosphere problem is usually considered as the determination of spatial dependence of basic macroscopic quantities \\citep[see also][\\citetalias{ja1}]{ja1} also be understood \\citep[following][]{ivandis} as aiming at finding three basic microscopic distributions, namely, the momentum distribution (distribution of velocities of all particles), distribution of particle internal degrees of freedom (populations of atomic excitation stages), and distribution of internal degrees of freedom of the electromagnetic field (radiation field for all frequencies and directions). ",
        "conclusions": "In this review we summarized the basic equations for the solution of the NLTE line formation problem with a focus on the solution of the problem for trace elements in stellar atmospheres. It has to be emphasized that in NLTE calculations for trace elements preference should be given to background NLTE model atmospheres, since using LTE model atmospheres is inconsistent with the condition of NLTE. It must be carefully decided which elements can be treated as a trace element and for which such treatment is questionable. After obtaining the numerical results all assumptions have to be verified, especially that of the traceness."
    },
    "1005/1005.0388_arXiv.txt": {
        "abstract": "Ongoing and upcoming surveys in x-rays and SZE are expected to jointly detect many clusters due to the large overlap in sky coverage. We show that, these clusters can be used as an ensemble of rulers to estimate the angular diameter distance, \\da. This comes at no extra observational cost, as these clusters form a subset of a much larger sample, assembled to build cluster number counts \\dn. On using this \\da, the dark energy constraints can be improved by factors of 1.5 - 4, over those from just \\dn. Even in the presence of a mass follow-up of 100 clusters (done for mass calibration), the dark energy constraints can be further tightened by factors of 2 - 3 . Adding \\da from clusters is similar to adding $d_{\\rm L}(z)$, from the SNe observations; for eg., \\dn (from ACT/SPT) plus \\da is comparable to \\dn plus $d_{\\rm L}(z)$ in constraining $\\Omega_m$ and $\\sigma_8$. ",
        "introduction": "Large cluster surveys like the SPT, ACT, Planck and eROSITA promise to detect from a thousand to a few hundred thousand clusters in the coming decade. The abundance and redshift distribution $\\frac{dN}{dz}$ of these clusters are important probes to understand the nature of dark energy as well as to constrain other cosmological parameters like $\\Omega_m$ and $\\sigma_8$ \\citep{Holder01, WangSteinhardt, Levine, Weller02, Hu}. To deduce a cosmology from these $\\frac{dN}{dz}$ observations, one requires a precise knowledge of the limiting mass of the survey as a function of redshift. One frequently uses proxy observables such as X-Ray surface brightness and temperature \\citep{Ebeling00}, Sunyaev Zel'dovich effect (SZE) decrement \\citep{Staniszewski, Hinks}, cluster richness \\citep{Postman, Koester} and lensing \\citep{Wittman06, Zitrin09} for the masses of clusters, related through simple power-law scaling relations. These scaling parameters are highly degenerate with the cosmological parameters, and breaking this degeneracy is crucial to obtain tight constraints on cosmology. This may realized, for example, through the so called `self-calibration' techniques  \\citep{MM04, Hu, LimaHu04}. Other approaches include an `unbiased' mass follow-up of a sub-sample of the survey clusters \\citep{MM03,MM04} or better theoretical modeling of clusters to predict the form of mass-observable scaling relation \\citep{YoungerHaiman, Reid06, CM09}. One can also try to optimize the cluster surveys so as to get the best possible survey yield \\citep{SatejPRL, Battye}.  Measurement of the angular diameter distance, $d_{\\rm A}$, at the redshift of the cluster using a combination of SZE and X-Ray observations have been routinely made over the last 30 years. The results have suffered, in the past, from various systematics and reliable estimates have only been achieved recently with analysis of statistically significant samples of galaxy clusters \\citep{Reese02,Bonamente}. These new observations have demonstrated the power of using clusters to measure `$d_{\\rm A}$ vs $z$' and use it to study the expansion history of our Universe \\citep{Molnar}. However, these recent progress has been done with targeted observations. Since, targeted observations are costly, this approach limits the size of the sample of $\\lsim 100$. In this {\\it letter}, we show that one can build up an `ensemble' of \\da by picking a sub-sample of clusters discovered in both X-Ray and SZE surveys with overlapping area and redshift coverage. Since, the surveys are already geared towards getting clusters for $\\frac{dN}{dz}$, we get the \\da without any {\\it extra} targeted observations. Addition of \\da to \\dn helps in tightening cosmological constraints, especially on dark energy equation of state. This is not surprising, since using \\da from clusters is akin to adding $d_{\\rm L}(z)$ information from supernovae (SNe) observations.  The rest of this paper is organized as follows. In \\S~2, we briefly discuss cluster number counts and the procedure to estimate \\da from SZE + X-Ray observations; in \\S~3 we describe the surveys, the choice of cosmological and cluster models and also summarize our methodology;  in \\S~4, we forecast constraints on cosmological parameters; and finally, we conclude in \\S~5. ",
        "conclusions": "We show that just using number count observations from the ongoing SZE galaxy cluster surveys like ACT/SPT would not be sufficient to put tight constraints on the cosmological parameters, especially the dark energy parameters. However these constraints can be significantly improved by doing a joint analysis with \\da. These $d_{\\rm A}(z)$ can be constructed out of clusters detected jointly in X-Ray and SZE observations having overlapping areas. The optical follow-up providing the redshifts for \\dn will, also, naturally provide the redshifts for \\da. With the current and upcoming large area surveys one will be able to get thousands of clusters providing us with us with $d_{\\rm A}$ at various redshifts without much effort.  We find that adding \\da to number count observations always improves the dark energy constraints, from \\dn alone by {\\it factors of 1.5 to 4}. This leads to better constraints not only on dark energy but also on the parameter $\\Omega_m$. Even when a targeted mass follow-up of clusters helps in breaking the cluster-cosmology degeneracies, addition of \\da helps in further tightening of the cosmological constraints. Moreover addition of \\da improves dark energy constraints comparable to the improvement brought by adding CMB priors on $\\sigma_8$ and $\\Omega_m$. Thus, our proposal of adding the \\da data to cluster number counts provides a natural way of improving the cosmological constraints using clusters alone.  The authors would like to thank Gil Holder, Jonathan Dudley and Joe Mohr for many discussions during the work."
    },
    "1005/1005.2344_arXiv.txt": {
        "abstract": "Nuclear hard X-ray luminosities (\\ln\\ba) for a sample of 112 early type galaxies within a distance of $67$ Mpc are used to investigate their relationship with the central galactic black hole mass \\mb (coming from direct dynamical studies or the \\mb$-\\sigma$ relation), the inner galactic structure (using the parameters describing its cuspiness), the hot gas content and the core radio luminosity. For this sample, \\ln ranges from $10^{38}$ to $10^{42}$ erg s$^{-1}$, and the Eddington ratio \\ln\\ba/\\ledd from $10^{-9}$ to $10^{-4}$, with the largest values belonging to four Seyfert galaxies. Together with a trend for \\ln to increase on average with the galactic luminosity \\lb and \\mb\\ba, there is a wide variation of \\ln (and \\ln\\ba/\\ledd\\ba), by up to 4 orders of magnitude, at any fixed \\lb$> 6\\times 10^{9}L_{B,\\odot}$ or \\mb$>10^7M_{\\odot}$.  This large observed range should reflect a large variation of the mass accretion rate \\md, and possible reasons for this difference are searched for. On the circumnuclear scale, in a scenario where accretion is (quasi) steady, \\md$\\,$ at fixed \\lb (or \\mb\\ba) could vary due to differences in the fuel production rate from stellar mass return linked to the inner galactic structure; a trend of \\ln with cuspiness is not observed, though, while a tendency for \\ln\\ba/\\ledd to be larger in cuspier galaxies is present.  In fact, \\md$\\,$ is predicted to vary with cuspiness by a factor exceeding a few only in hot gas poor galaxies and for large differences in the core radius; for a subsample with these characteristics the expected effect seems to be present in the observed \\ln values. \\ln does not show a dependence on the age of the stellar population in the central galactic region, for ages$>$3 Gyr; less luminous nuclei, though, are found among the youngest galaxies or galaxies with a younger stellar component.  On the global galactic scale, \\ln shows a trend with the total galactic hot gas cooling rate ($L_{X,ISM}$): it is detected both in gas poor and gas rich galaxies, and on average increases with $L_{X,ISM}$, but again with a large scatter.  The observed lack of a tight relationship between \\ln and the circumnuclear and total gas content can be explained if accretion is regulated by factors overcoming the importance of fuel availability, as 1) the gas is heated by black hole feedback and \\md$\\,$ varies due to an activity cycle, and 2) the mass effectively accreted by the black hole can be largely reduced with respect to that entering the circumnuclear region, as in radiatively inefficient accretion with winds/outflows.  Finally, differently from \\ln\\ba, the central 5 GHz VLA luminosity shows a clear trend with the inner galactic structure, that is similar to that shown by the total soft X-ray emission; therefore it is suggested that they could both be produced by the hot gas. ",
        "introduction": "In the past years, high angular resolution studies of the centers of early type galaxies have been performed with the $Hubble$ $Space$ $Telescope$ ($HST$) in the optical and near infrared, and in the X-rays with $Chandra$, obtaining important results that deeply influenced our understanding of the nature and past evolution of these systems.  The first, major $HST$ result was that massive black holes (MBHs) are ubiquitous in the centers of spheroids, and linked by tight relationships with the luminosity and central stellar velocity dispersion of their hosts (e.g., Magorrian et al. 1998, Ferrarese \\& Merritt 2000, Gebhardt et al. 2000), indicative of a strong mutual influence during their formation and evolution. The second $HST$ result was that the central brightness profiles of galaxies with $M_V\\lsim -19$ show either steep brightness cusps or, interior to a break radius $r_b$, they flatten markedly in a core with respect to an inner extrapolation of the outer profile.  These profiles have been described respectively by a \\Ser or core-\\Ser law (Graham et al. 2003, Trujillo et al. 2004, Ferrarese et al. 2006, Kormendy et al. 2009) or alternatively by the \"Nuker law\" (Faber et al. 1997, Lauer et al. 2007a). Cores dominate at the highest luminosities and steep cusps at the lowest, with an intermediate luminosity region of coexistence ($-20.5\\gsim M_V \\gsim -23$, Lauer et al. 2007a).  The shape of the brightness profile in the inner galactic region has been related to the past formation and evolution of galaxies, with cores created during dry merging events by a black hole binary ejecting stars from the center of the new system (Ebisuzaki et al. 1991, Faber et al. 1997, Milosavljevic et al. 2002, Graham 2004, Gualandris \\& Merritt 2008), and cusps being preserved or (re)generated during gaseous (wet) mergings.  Recently, it was found that coreless elliptical galaxies in the Virgo cluster have extra-light at their center, above the inward extrapolation of their outer \\Ser profile (Kormendy et al. 2009), the result of a wet merger induced starburst (see also Hopkins et al. 2009a) or of AGN induced starburst activity (Ciotti \\& Ostriker 2007).  Moreover, the presence of steep cusps or cores correlates with other fundamental galactic properties, even more tightly than how these properties correlate with the galatic luminosity: core galaxies generally have boxy isophotes, are slow rotators and triaxial systems, while cusp galaxies are disky, fast rotators and axisymmetric (Kormendy \\& Bender 1996, Faber et al. 1997); core galaxies show a large range of radio and X-ray luminosities, while cusp galaxies are confined below a threshold (Bender et al. 1989; Pellegrini 1999, 2005a; Capetti \\& Balmaverde 2005; Pasquali et al. 2007).  Pellegrini (2005a) also attempted an investigation of the relation between the X-ray nuclear emission (\\ln) and the inner core/cusp profile, but the study was limited by the small number of nuclei with known \\ln available.  Launched in 1999, the $Chandra$ satellite has now pointed a large number of early type galaxies, with an unprecedented angular resolution in the X-rays of less than $1^{\\prime\\prime}$.  For the first time measurements of the nuclear X-ray emission down to values as low as $10^{39}$ erg s$^{-1}$ and out to distances of $\\sim 60$ Mpc have been obtained.  The MBHs of the local universe turned out to be typically radiatively quiescent and very sub-Eddington emitters (Loewenstein et al. 2001, Soria et al. 2006a, Zhang et al. 2009, Gallo et al. 2008, 2010). In a number of cases the mass accretion rate on the MBH could be estimated (e.g., Di Matteo et al. 2003, Pellegrini 2005b) and the radiative quiescence was interpreted in terms of radiatively inefficient accretion (Narayan \\& Yi 1995), possibly with the mechanical power dominating the total output of accretion (e.g., Allen et al. 2006).  However, many aspects of accretion in the local universe remain unknown: what determines \\ln ?  Is there any relation of \\ln or its Eddington-scaled value with the galactic luminosity or the mass of the central supermassive black hole \\mb\\ba?  or with the inner stellar profile, that has been linked to the past galactic evolution and other major global galactic properties? or with the hot gas content?  Answering these questions is important for a complete understanding of the MBH-host galaxy coevolution process.  In this work we have collected all early type galaxies (E and S0) out to a distance of $\\sim 67$ Mpc with known \\ln\\ba, based mostly on data coming from $Chandra$ pointings; a total of 112 galaxies resulted with \\ln measured or with an upper limit on it.  The sample includes also most of the early type galaxies with a direct measurement of \\mb via dynamical studies currently available; for the other galaxies, the central stellar velocity dispersion allows for an estimate of \\mb via the \\mb$-\\sigma$ relation.  For 81 of these 112 galaxies the central stellar profile shape has been measured with $HST$.  The sample is described in Sect.~\\ref{sample}, the observational evidences about relationships between \\ln\\ba, \\mb\\ba, the central stellar structure and the radio luminosity are presented in Sect.~\\ref{obs}, the results are discussed in Sect.~\\ref{disc} (that examines also the relationship between \\ln and the galactic hot gas luminosity); the conclusions are summarized in Sect.~\\ref{concl}. ",
        "conclusions": "In this work measurements or upper limits have been collected for the hard X-ray emission at the nucleus of 112 early type galaxies (E and S0) of the local universe (within 67 Mpc). \\ln derives from $Chandra$ data for 94\\% of the nuclei, and the sample includes all the available measurements of \\ln (29 cases) for galaxies with the above characteristics and a direct estimate of the \\mb (38 objects). Using also \\mb from the \\mb$-\\sigma$ relation, and the inner stellar profile (slope $\\gamma$ and break radius $r_b$) measured with $HST$ for 81 galaxies, the relationships between \\ln and \\mb\\ba, the central stellar structure, the central age, the radio and the soft X-rays (hot gas) luminosities have been investigated, with the following results:  $\\bullet$ \\ln increases on average with \\lb and \\mb\\ba, with a wide variation of \\ln and \\ln\\ba/\\ledd\\ba, up to 4 orders of magnitude, at any fixed galactic \\lb$> 6\\times 10^{9}L_{B,\\odot}$ or \\mb$>10^7M_{\\odot}$. Cusp and core galaxies both cover the whole large range of \\ln\\ba, without a clear trend with $\\gamma$ or $r_b$.  Most nuclei have an Eddington ratio $-9<$log(\\ln\\ba/\\ledd\\ba)$<-5$, the brightest ones (with log(\\ln\\ba/\\ledd\\ba)$\\sim -4$) are four Seyfert nuclei. The  \\ln\\ba/\\ledd is highest at the lowest \\mb\\ba, and shows an increase of its highest values with increasing \"cuspiness\" ($\\gamma$).  $\\bullet$ Accretion in these MBHs may have entered the radiatively inefficient regime, where \\ln scales as the mass accretion rate \\md$^2$ at fixed \\mb\\ba; therefore reasons for \\md$\\,$ variations explaining the wide range of observed \\ln at fixed \\lb (or \\mb\\ba) are searched for.  In a scenario where accretion is (quasi) steady, \\md$\\,$ could vary because of a different inner galactic structure, that determines the amount of mass shed by stars in the circumnuclear region and then available for accretion. The lack of a clear trend of \\ln with $\\gamma$ and $r_b$, and the weak evidence for an increase of \\ln\\ba/\\ledd with cuspiness, indicate that differences in the inner galactic structure do not have the dominant effect on \\ln\\ba. In fact the circumnuclear fuel production rate is estimated to vary more than a factor of a few only when the inflowing region is small (of the order of the accretion radius) and for large differences in $r_b$; in agreement with this, for galaxies with presumably a small inflowing region, a trend for \\ln to decrease for increasing $r_b$ seems to be present.  $\\bullet$ The stellar mass return rate depends also on the age of the stellar population, but \\ln covers the same wide range of values at all central ages from 3 to 14 Gyr; \"youth\", indicated by a younger center, recent starformation, and/or a younger stellar component at the center (of ages $<2-3$ Gyr), seems to be more connected with a lower \\ln\\ba. An explanation could be that the heating by starformation or the dynamical effect of the merging process on the hot gas drive the available gas out in a wind, ending abruptly the starformation and accretion.  However, this finding remains an indication, because of the trend of \\ln with \\lb and the small number of galaxies with information about recent starformation episodes.  $\\bullet$ \\ln on average increases with the total hot gas emission $L_{X,ISM}$, but the relation is not tight: it shows a large variation of 2--3 orders of magnitude both at the lowest and highest gas contents.  Nuclear emission at the highest detected levels for this sample (\\ln$\\sim 10^{40}-10^{41}$ erg s$^{-1}$) is present even when the gas content is low, and \\ln at the lowest levels ($\\sim 10^{38}-10^{39}$ erg s$^{-1}$) can be found at gas luminosities differing by 3 orders of magnitude.  $\\bullet $ The mild sensitivity of \\ln to the circumnuclear and global hot gas contents, and its large variation, finds two possible explanations, both of which go in the sense of overcoming the importance of fuel availability : 1) the gas is heated due to feedback from the MBH, which could be a phenomenon not limited to high $L_{X,ISM}$ galaxies, or 2) only a small fraction of the mass entering the accretion radius actually reaches the MBH, as in RIAF models with winds/outflows, and this fraction can vary largely at fixed \\mb\\ba.  $\\bullet $ A sub-sample of galaxies shows an already known trend by which cusp galaxies are confined below a threshold in the VLA 5 GHz central luminosity $L_{R,core}$, while core galaxies span a large range of $L_{R,core}$; this subsample does not show a similar trend of \\ln with the central light profile.  The $L_{R,core}$ behavior is instead similar to that of the total soft X-ray emission with respect to $\\gamma$; the hot gas could then be responsible both for the soft X-ray emission and the jet confinement and propagation. While core galaxies can possess a conspicuous radio activity cycle, in cusp galaxies the variation of the radio emission keeps within a smaller range, because of a rapid jet failure due to the lack of a dense confining medium, or a smaller duty cycle, being these galaxies on average less massive systems."
    },
    "1005/1005.1171_arXiv.txt": {
        "abstract": "% The effect of massive neutrinos on matter power spectrum is discussed in the context of $f(R)$ gravity. It is shown that the anomalous growth of density fluctuations on small scales due to the scalaron force can be compensated by free streaming of neutrinos. As a result, models which predict observable deviation of the equation-of-state parameter $w_\\DE$ from $w_\\DE=-1$ can be reconciled with observations of matter clustering if the total neutrino mass is $O(0.5~\\eV)$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.1940_arXiv.txt": {
        "abstract": "When studying the evolutionary stages of protostars that form in clusters, the role of any intracluster medium cannot be neglected. High foreground extinction can lead to situations where young stellar objects (YSOs) appear to be in earlier evolutionary stages than they actually are, particularly when using simple criteria like spectral indices. To address this issue, we have assembled detailed SED characterizations of a sample of 56 \\textsl{Spitzer}-identified candidate YSOs in the clusters NGC~2264 and IC~348. For these, we use spectra obtained with the Infrared Spectrograph onboard the \\textsl{Spitzer} Space Telescope and ancillary multi-wavelength photometry. The primary aim is twofold: 1) to discuss the role of spectral features, particularly those due to ices and silicates, in determining a YSO's evolutionary stage, and 2) to perform comprehensive modeling of spectral energy distributions (SEDs) enhanced by the IRS data. The SEDs consist of ancillary optical-to-submillimeter multi-wavelength data as well as an accurate description of the 9.7~$\\mu$m silicate feature and of the mid-infrared continuum derived from line-free parts of the IRS spectra. We find that using this approach, we can distinguish genuine protostars in the cluster from T Tauri stars masquerading as protostars due to external foreground extinction. Our results underline the importance of photometric data in the far-infrared/submillimeter wavelength range, at sufficiently high angular resolution to more accurately classify cluster members. Such observations are becoming possible now with the advent of the \\textsl{Herschel} Space Observatory. ",
        "introduction": "Finding and characterizing individual Young Stellar Objects (YSOs) in embedded clustered environments is challenging because the observed YSO spectral energy distributions (SEDs) can easily be influenced by dense material that is distributed throughout the cluster, but physically unrelated to an individual YSO or its envelope. In that situation, the extra extinction can mislead the observer since an affected source may appear to be in an earlier evolutionary stage than it actually is. For example, the SED of a T Tauri star behind foreground extinction can in some circumstances look protostellar. In a more isolated environment, this situation would only occur in case of edge-on disks, aggravated by the presence of disk flaring. To avoid confusion, \\citet{rob06} proposed that sources be assigned designations that refer to the actual evolutionary stages, rather than spectral classes defined by observations, since the latter are affected by extinction. In most instances, the evolutionary stages (I, II, III) correspond to the SED classes. However, in heavily extincted regions the SED classes may not always correspond to the appropriate evolutionary stages. In embedded clustered environments, the difference between YSO classes and evolutionary stages becomes particularly important.  Infrared spectroscopy allows us to go beyond protostellar classifications based on photometry. At near-infrared wavelengths, varying degrees of veiling can be used to distinguish protostars from PMS stars \\citep{cas92,gre96,whi07}. At mid-infrared wavelengths, silicate dust and ices cause telltale spectral features that can be used to refine spectral classes. Many of these features occur in the wavelength range that was accessible with the Infrared Spectrograph (IRS) onboard the \\textsl{Spitzer} Space Telescope \\citep{hou04}. In the context of star formation, the power of this instrument has been shown most notably by studies of T Tauri stars and protostars in Taurus, Chamaeleon, and Ophiuchus \\citep{fur06,fur08,fur09}. Most recently, \\citet{mcc10} study an extensive sample of YSOs in Ophiuchus, using IRS data to determine their evolutionary stages.  To find out whether the effect of intracluster or external foreground extinction can be disentangled from the actual evolutionary stage of a YSO, we have obtained IRS spectra and assembled detailed multi-wavelength photometry for a sample of 56 candidate YSOs in NGC~2264 and IC~348. We then use these data for comprehensive SED modeling with the primary aim of ascertaining the evolutionary stages for the sources in our sample. These can then finally be compared to previous classification attempts that are mainly based on less broad mid-infrared wavelength coverage. For each of the 56 sources, we have obtained mid-infrared spectra using \\textsl{Spitzer}-IRS. These data not only allow us to study spectral features that trace the evolutionary stage of a YSO (silicates and ices), but they also allow us to derive photometry from line-free parts of the spectrum to accurately reflect the underlying continuum. This technique also allows us to accurately describe the main silicate spectral feature by photometric data at selected wavelengths. Combined with ancillary photometric data spanning optical to submillimeter wavelengths, the IRS data thereby allow us to construct SEDs that are tailored toward the requirements of detailed modeling. The accuracy of the modeling depends on the SEDs not containing spectral lines that are not part of the underlying model (e.g., PAH features); these features can fall into the band passes of broadband filters. Our IRS-derived photometry attempts to depict the continuum and the silicate features. Compared to the usual approach of using photometry from \\textsl{Spitzer}-IRAC in SEDs for subsequent modeling, we discard bands 3 and 4 of the IRAC data that we have for every source in favor of corresponding line-free photometry derived from the IRS data.   \\subsection{Source selection}  We selected 56 candidate protostars in NGC 2264 and IC\\,348 -- 40 in the former and 16 in the latter region -- from previous \\textsl{Spitzer} photometric studies of these regions by \\citet{mue07} and \\citet{tei08}. Our target sources are listed in Table~\\ref{tab_srclist}.  NGC 2264 is a prominent, clustered star-forming region located at a distance of $\\sim$913~pc \\citep{bax09}; for a recent review, see \\citet{dah08}. Here, we focus on the population of embedded protostars in the vicinity of IRS~2 \\citep{you06} where we selected 40 sources for \\textsl{Spitzer}-IRS mid-infrared spectroscopic observations. This region is one of the two most prominent regions of ongoing star formation in NGC~2264, next to IRS~1. All selected sources are located in the \\textsl{Spokes} cluster \\citep{tei06}. Figure~\\ref{mips_n2264} shows the selected sources in a mid-infrared view of the region. IC\\,348 is located in Perseus at a distance of $\\sim$260~pc \\citep[e.g.,][]{lom10}; for a recent review, see \\citet{her08}. \\citet{mue07} conducted a \\textsl{Spitzer} census of this star-forming region. The IC\\,348 sources discussed in this paper are among the most embedded in the cluster and are predominantly protostars in a narrow filamentary ridge located about 1~pc southwest of the central cluster. In Figure~\\ref{mips_ic348}, we mark the selected sources in a mid-infrared view of the region.  A first comparison within our sample can be accomplished by using the IRAC spectral index $\\alpha_{\\rm IRAC}$, as determined from a least-squares fit to the photometry of all four IRAC bands \\citep{lad06}; see section~\\ref{sec_anc} for a discussion of the photometry. The sources in our NGC 2264 subsample show diverse IRAC spectral indices, ranging from $\\alpha=-1.41$ (source 27) to $\\alpha=4.21$ (source 45). Eighteen sources in this subsample have positive spectral indices. \\citet{sun09} identify 37 out of the 40 sources as candidate cluster members in various stages, 20 of them are listed as class~I sources; their classification is included in Table~\\ref{tab_srclist}. Note that while \\citet{sun09} list no membership status for the three remaining sources, our results suggest that they are cluster members, too. In the IC\\,348 sample, the IRAC spectral indices range from $\\alpha=-0.93$ (source 14) to $\\alpha=1.92$ (source 6). The classification from \\citet{lad06} and \\citet{mue07} is listed in Table~\\ref{tab_srclist}. Most of the sources in this subsample are class I objects. Since we do not find significant differences between these two cluster subsamples -- both (unevenly) sample the same YSO classes --,  we will discuss them as one sample in the following, focusing more on a discussion of evolutionary stages in general. Note, however, that the source table allows the reader to identify individual YSOs in both clusters.  In this paper, we first discuss the \\textsl{Spitzer}-IRS observations and their data reduction. We then describe the ancillary multi-wavelength data, ranging from optical to submillimeter wavelengths, and how we assembled composite SEDs from both IRS spectra and ancillary data. Subsequently, we perform comprehensive modeling of these SEDs before concluding with a discussion and summary of the results. ",
        "conclusions": "After the analysis of the IRS spectra and subsequent detailed SED modeling, we can characterize the YSOs in our sample more accurately than what was previously possible from YSO spectral classes alone by assigning evolutionary stages next to their SED classes. As we have seen, the comprehensive SED modeling suggests that the sources fall into three classes: 1) genuine protostars (stage I), 2) disk sources behind high foreground extinction (stage II(ex)), and 3) disk sources with less external foreground extinction (stage II). Of the 44 sources for which we determine evolutionary stages, 17 are stage I protostars. Most of these were already classified as class~I in the literature (one, source 4,  is a class 0 source, but a stage earlier than stage~I has not been defined), with the exceptions of sources 28, 45, and 51 which have not been classified, or have been classified differently (source 35) in \\citet{sun09}. On the other hand, as summarized in Table~\\ref{tab_srclist}, the sample contains many previous class~I sources that instead appear to be in stage II(ex). Thus, a number of sources appear to be in later evolutionary stages than what was previously suspected. Most previously determined class~II sources turn out to be in stage II, although a few are in stage II(ex).  With regard to the use of spectral indices as an indicator of YSO class, our sample contains a few objects that can serve as instructive examples of the limits of using a single spectral index for the classification of individual sources. Sources 3 and 56 both have negative IRAC spectral indices, but it turns out that they are stage I protostars and were classified as class~I protostars in the literature.  The IRS data are a significant help to constrain the SEDs of these sources both by allowing us to make use of spectral properties in YSO characterization, and by extending the SED wavelength coverage. The latter is further improved since the IRS data allow us to avoid spectral ranges contaminated by unwanted line features for photometry and subsequent SED modeling. In particular, the youngest sources show strong silicate and ice features in their mid-infrared spectra; both features are correlated with the sources' spectral indices, both IRAC and \\textsl{K}--24\\,$\\mu$m. A correlation of the silicate optical depth at 9.7\\,$\\mu$m and the total ice column density is present, but has more scatter than the correlations of the two features with the IRAC spectral index as a measure of evolutionary stage. The main exception appear to be objects with high silicate optical depths but low total ice column densities. These might be more massive stars with radiation fields able to evaporate the ice mantles. A similar correlation is found in the c2d data and by \\citet{fur08}. We detect three sources with PAH emission (all in NGC~2264), a small subset of our sample. Presumably, these are more massive YSOs. For the crowded region of the \\textsl{microcluster} of NGC~2264, we can make use of IRS spectral maps, for example to identify regions with jet activity.  Based on the IRS and modeling results, we can check whether the three evolutionary stages that were identified by SED modeling in turn have distinctive features in their IRS spectra. We have therefore marked the sources in Fig.~\\ref{fig_speccorr} according to their evolutionary stage. Indeed, the stage I sources tend to have large optical depths of the silicate feature and strong ice features. Stage~II(ex) sources have less strong silicate and ice features and stage II sources tend to show the silicate feature in emission while not showing any ice features. It is interesting to note that the IRS spectral features also appear to be distinctly different when comparing stage II(ex) and stage II sources. Fig.~\\ref{fig_speccorr} shows that while virtually all stage II sources have the 9.7~$\\mu$m in \\textsl{emission}, indicative of disks, that is \\textsl{not} the case for the stage II(ex) sources, presumably due to the additional extinction which may `cancel out' the emission. Also, none of the stage II sources show ice features, while stage II(ex) interestingly do show ice absorption, although at smaller total column densities than the stage I sources. In our interpretation, the stage II(ex) ice features would be related to the additional extinction by intracluster foreground material and not the sources themselves. In summary, while IRS data are of crucial importance, they alone are not sufficient to distinguish the evolutionary stages. For this purpose, detailed SED modeling is required.  We point out that in the context of our comprehensive SED modeling, the archival submillimeter data were very important for the separation of the sample into the different evolutionary stages. Dense dust and gas in an envelope manifest themselves in submillimeter maps as concentrated emission in contrast to more large-scale foreground extinction. Future high-angular resolution observations in the far-infrared/submillimeter wavelength range, as forthcoming from the \\textsl{Herschel} Space Observatory, will be even better suited to help determining the evolutionary stages in YSO samples.  Finally, we discuss the observable quantities that correlate best with the different evolutionary stages. We have seen that this is certainly not a spectral index since the correlation of the evolutionary stages with spectral indices is not unambiguous. The genuine protostars (stage I) all have submillimeter detections and the highest silicate optical depths as well as the highest ice column densities. The stage II sources are broadly characterized by a silicate feature in emission and no ice features. The stage II(ex) sources lie between these two extremes, but their SEDs can be easily mistaken for class I sources (in fact, some were previously determined to be class I sources, see above); in particular, they do not show the silicate feature in emission. Once an individual SED shape is well characterized, far-infrared and submillimeter continuum data are the most important photometric data to distinguish the different evolutionary stages."
    },
    "1005/1005.1037_arXiv.txt": {
        "abstract": "We continue to explore the validity of the reflected shock structure (RSS) picture in \\snr that was proposed in our previous analyses of the X-ray emission from this object. We used an improved version of our RSS model in a global analysis of 14 CCD spectra from the monitoring program with {\\it Chandra}. In the framework of the RSS picture, we are able to match both the expansion velocity curve deduced from the analysis of the X-ray images  and light curve. Using a simplified analysis, we also show that the X-rays and the non-thermal radio emission may originate from the same shock structure (the blast wave). We believe that using the RSS model in the analysis of grating data from the {\\it Chandra} monitoring program  of \\snr that cover a long enough time interval, will allow us to build a more realistic physical picture and model of \\snrE. ",
        "introduction": "After its reappearance in radio (\\citealt{turtle_90}; \\citealt{stav_92}, 1993) and in X-rays (\\citealt{buer_94}; \\citealt{gor_94}; \\citealt{has_96}) at about 1000-1200 days after the explosion (DAE), SN 1987A entered into the phase of supernova remnant (SNR) formation \\citep{dick_07}. Observations of this phenomenon from its onset give us a chance to study in detail the physics of strong shocks as well as to investigate the distribution of the circumstellar matter (CSM) around the exploded star. The latter may help us solve the mystery of the origin of the triple-ring system observed in the optical (\\citealt{wampler_90}; \\citealt{jak_91}; \\citealt{crotts_91}; \\citealt{burr_95}). With this respect, X-ray observations are very important since they manifest most of the energetics at present and they provide us with direct information about the underlying physics of this evolving object.  The richness and variety of processes involved in the birth of \\snr are manifested by observations across the entire electromagnetic spectrum. Radio observations (\\citealt{stav_92}, 1993, 2007; \\citealt{ball_01}; \\citealt{man_02}, 2005; \\citealt{gaensler_97}, 2007; \\citealt{ng_08}; \\citealt{potter_09}; \\citealt{zana_10}) tell us about the non-thermal (synchrotron) emission mechanisms and relativistic particle acceleration in strong shocks. Near-infrared observations tell us about the distribution of dust in the CSM and the physics of its interaction with the hot plasma (\\citealt{bou_04}, 2006; \\citealt{dwek_08}). Optical and UV observations (\\citealt{son_98}; \\citealt{law_00}; \\citealt{pun_02}; \\citealt{suger_02}; \\citealt{gron_06}, 2008) tell us about the physics of slow shocks and the distribution of dense clumps of gas.   The early evolution of \\snr is strongly influenced by the distribution of circumstellar gas, which is concentrated in its triple-ring system. Although the origin of the ring system is not well understood, the colliding wind scenario that assumes a spherically-symmetric blue supergiant (BSG) wind interacting with an earlier emitted asymmetric red supergiant (RSG) wind seems the physically most reasonable explanation (\\citealt{luo_91}; \\citealt{wang_92}; \\citealt{morris_07}). It has been explored in detail by \\citet{bl_lu_93} who demonstrated that such a model is capable of accounting for the observed CSM characteristics, provided the RSG wind was highly asymmetric with most of its mass confined near the equator. This picture needed an additional ingredient in order to explain the reappearance of \\snr in radio and X-rays. \\citet{che_dwa_95} proposed that an extended HII region was formed inside the inner equatorial ring that was naturally produced by the photoionization of the shocked RSG wind (the inner ring) by the central progenitor BSG star. Based on this, Borkowski et al. (1997a, b) studied the complex hydrodynamics of the interaction of the SN ejecta with the CSM in vicinity of the inner ring that also predicts a continuous brightening of \\snr in UV and X-rays.  Since its reappearance, \\snr has been brightening continuously in X-rays with an accelerating rate in the past few years (\\citealt{sang_05}, 2006, 2007; \\citealt{judy_09}). The first spectrum of \\snr with high spectral resolution was obtained with the high energy transmission grating on {\\it Chandra} by \\citet{eli_02}. The subsequent  brightening has enabled dispersed spectra with good photon statistics to be obtained with both {\\it Chandra} (\\citealt{zhek_05}, 2006; \\citealt{dew_08}; \\citealt{zhek_09}) and {\\it XMM-Newton} (\\citealt{haberl_06}; \\citealt{heng_08}). With the superb spatial and spectral resolution of the {\\it Chandra} observatory, the first direct evidence was obtained that the distribution of the X-ray emitting plasma in \\snr is not spherically-symmetric but it is `flat', {\\it i.e.,} it is concentrated in the vicinity of the equatorial ring (\\citealt{zhek_05}, 2006; \\citealt{dew_08}; \\citealt{zhek_09}). These authors also showed that the bulk gas velocities deduced from the X-ray spectral lines are too low to account for the plasma temperatures inferred from the spectral fit, and showed that a simplified reflected-shock structure (RSS) model could match the observed spectra successfully.   Superb spatial resolution of {\\it Chandra} allowed \\citet{burr_00} to obtain the first  image of \\snr that showed a ring-like structure in X-rays. Subsequently, a monitoring program was established with {\\it Chandra} to track the evolution of the X-ray images and CCD spectra of \\snrE. In addition to  the continuing brightening in X-rays, the spectral data showed a gradual decrease of the plasma temperature  behind the fast shocks. From the imaging data, we measured the expansion velocity of the remnant (\\citealt{sang_02}, 2004, 2007).  After applying a physically more realistic method to analyze the X-ray images, \\citet{judy_09} obtained a more accurate expansion velocity curve. It showed clearly that the velocity started to decelerate after $\\sim 6000$~DAE, about the same time that the X-ray light curve turned up \\citep{sang_05}.  Here we analyze the evolution of the X-ray spectrum of \\snr as observed with the pulse-height spectra obtained with the ACIS CCD camera during the monitoring program.  One of the main goals of our study is to test the reflected-shock structure picture that emerged from analysis of the {\\it Chandra} grating data. Since the pulse-height spectra have much better time coverage than the grating data, they can test whether is possible to match the CCD spectra and the expansion velocity curve in the framework of the RSS picture. In section \\S~\\ref{sec:rss_picture} we give a qualitative description of the RSS picture for \\snrE. In section \\S~\\ref{sec:rss_model} we present details about the global RSS model. We describe the evolution of the CCD spectra and results from the spectral fits in sections \\S~\\ref{sec:data} and \\S~\\ref{sec:spectral_fits}, respectively. In section \\S~\\ref{sec:discussion} we discuss the results, how they might be related to the non-thermal radio emission from \\snr and the consistency of the global RSS picture. We present our conclusions in section \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} In this study, we continued to explore the validity of the reflected-shock structure picture in \\snrE. This picture assumes that the blast wave is interacting with the HII region in vicinity of the inner ring and when it encounters dense clump(s) a shock is transmitted into the clump(s) and a reflected shock forms that additionally compresses the shocked gas behind the blast wave. We adopted an improved version of the reflected-shock model (global RSS model) to test the RSS picture in \\snr by making use of CCD spectra from the monitoring program with {\\it Chandra}. The main results from this analysis are as follows.  1. The global RSS model is capable of matching both the X-ray expansion velocity curve \\citep{judy_09} as well as the X-ray spectra of \\snr over its evolution for the last ten years or so. Moreover, extrapolating back in time the X-ray flux, the model is able to match the observed {\\it ROSAT} fluxes from the time of the supernova reappearance in X-rays (1000-3000 DAE).  2. The evolution of the mean electron temperature behind the shocks shows a decrease  with time of the electron temperature in the blast wave, while a slight increase with time is noticed for this parameter in the transmitted shock. The mean ratio of the electron to the mean plasma temperature is required to increase with the decrease of the velocity of the blast wave.  3. At early times, the soft (0.5-2 keV) X-ray flux of \\snr was dominated by the blast wave but for times after $\\sim 6000$ DAE the transmitted shock became dominant. This transition coincides with the steep upturn in the X-ray LC reported by \\citet{sang_05}. Also, the emission measure of the hot plasma in the transmitted shock (as well as in the reflected shock) is increasing steeper than that for the blast wave for times after $\\sim 6400$ DAE. All this is indicative of an ongoing penetration of the blast wave deeper in the equatorial ring debris.  4. Results from the X-ray analysis were used to explore the possibility that the X-rays and the non-thermal radio emission from \\snr originate from the same shock structure (the blast wave). Using a simplified analysis, it was shown that this can indeed be the case, provided the particle acceleration mechanisms, that operate in the shocks in \\snrE, are not very efficient in producing relativistic particles with very high energies.  5. As a next step, the global RSS model could be used to obtain further constraints on the reflected-shock structure picture in \\snr by analyzing X-ray spectra with good spectral resolution once the available data cover a long enough time interval of the \\snr evolution. We believe that this will allow us to improve our understanding of the underlying physics and to build a more realistic picture and model of this fascinating phenomenon: the birth and evolution of supernova remnant 1987A."
    },
    "1005/1005.3817_arXiv.txt": {
        "abstract": "We report spectroscopic confirmation and high-resolution infrared imaging of a $z=2.79$ triply-imaged galaxy behind the Bullet Cluster. This source, a \\spitzer-selected luminous infrared galaxy (LIRG), is confirmed via polycyclic aromatic hydrocarbon (PAH) features using the \\spitzer~Infrared Spectrograph (\\irs) and resolved with \\hst~\\wfc~imaging. In this galaxy, which with a stellar mass $M_*\\approx4\\times10^9$ M$_\\sun$ is one of the two least massive ones studied with \\irs\\ at $z>2$, we also detect $H_2\\;S(4)$ and $H_2\\;S(5)$ pure rotational lines (at 3.1$\\sigma$ and 2.1$\\sigma$) -- the first detection of these molecular hydrogen lines in a high-redshift galaxy. From the molecular hydrogen lines we infer an excitation temperature $T=377^{+68}_{-84}$ K. The detection of these lines indicates that the warm molecular gas mass is $6^{+36}_{-4}$\\% of the stellar mass and implies the likely existence of a substantial reservoir of cold molecular gas in the galaxy. Future spectral observations at longer wavelengths with facilities like the \\herschel\\ {\\it Space Observatory}, the Large Millimeter Telescope, and the Atacama Pathfinder EXperiment (APEX) thus hold the promise of precisely determining the total molecular gas mass. Given the redshift, and using refined astrometric positions from the high resolution imaging, we also update the magnification estimate and derived fundamental physical properties of this system. The previously published values for $L_{IR}$, star formation rate, and dust temperature are confirmed modulo the revised magnification; however we find that PAH emission is roughly a factor of five stronger than would be predicted by the relations between $L_{IR}$ and $L_{PAH}$ reported for SMGs and starbursts in \\citet{pope2008}. ",
        "introduction": "\\label{sec:intro}   An important legacy of the \\spitzer~{\\it Space Telescope} is that it enabled the first detection of dust emission from a large number of luminous infrared galaxies at $z>2$ \\citep[e.g.,][]{papovich2006,perez2005}, demonstrating that these galaxies dominate the massive galaxy population at high redshift. The implication is that at $z\\sim 2-3$ massive galaxies are rapidly assembling their stars and growing supermassive black holes.  For the most infrared-luminous systems, spectra from the \\spitzer\\ Infrared Spectrograph (IRS) have provided important insights into the physical processes driving $L_{IR}$ \\citep[e.g.,][]{houck2005,yan2007}. \\citet{pope2008} used \\irs\\ to demonstrate that the mid-IR properties of ultraluminous infrared galaxies (ULIRGs, $L_{IR}>10^{12}$ L$_\\sun$) and submillimeter galaxies (SMGs) are distinct, with star-formation dominating the infrared emission for typical SMGs.  It has been postulated by this group and others that such differences indicate different evolutionary stages, as the dominant source of $L_{IR}$ transitions from star formation to AGN emission.  Existing studies like \\citet{pope2008} clearly demonstrate that \\irs~ spectroscopy provides a clean means of disentangling the AGN and star formation contributions. A limitation to such work though is that current samples only probe the bright end of the luminosity function at high-z \\citep{dey2008,dye2008}, and \\irs~observations are practical for only a small subset of the most luminous sources. Observations of luminous infrared galaxies (LIRGS, $10^{11}<L_{IR}<10^{12}$ L$_\\sun$), the dominant class of IR galaxies at these redshifts \\citep[$\\sim10\\times$ higher density than ULIRGS]{caputi2007}, have simply not been feasible.  Gravitational lensing provides a means of circumventing this limitation, however, and several recent programs exploit strong lensing to study the properties of lensed galaxies at mid-IR and submillimeter wavelengths \\citep[][and references therein]{knudsen2008,rigby2008}.  In this paper we utilize a similar strategy to study a highly magnified luminous infrared galaxy that lies behind the Bullet Cluster \\citep{markevitch2002}.  First identified in \\spitzer~\\irac~imaging \\citep{bradac2006}, this galaxy has been shown to be a triply imaged optical drop-out \\citep[][hereafter \\firstp]{gonzalez2009} that is extremely luminous at submillimeter wavelengths \\citep[$\\sim 0.11$ Jy at 500$\\mu$m][]{wilson2008,blast2009,johansson2010,rex2010}. Photometric redshift estimates based upon both the stellar SED and the submillimeter emission place the galaxy at $z\\sim2.7-2.9$ \\citep{wilson2008,gonzalez2009,blast2009}, while lensing constraints imply magnification factors of $10-50$ for each of the three images \\firstpaper.  This galaxy therefore presents a unique opportunity to study in detail the properties of a member of the LIRG population during the peak epoch of star formation and galaxy assembly when these galaxies may be responsible for more than half the total infrared emission \\citep{rodighiero2010}.  We present in this paper the results of targeted observations of this lensed galaxy with \\spitzer~\\irs~spectroscopy and \\hst\\ Wide Field Camera 3 (\\wfc) infrared imaging. We combine these observations with both refined positions from the \\wfc\\ imaging and published results from \\citet{wilson2008}, \\firstp, and \\citet{blast2009} to derive updated physical parameters for this galaxy. We also constrain the total contribution of a central active galactic nuclei (AGN) to the total mid-infrared emission. The data are presented in \\S\\ref{sec:data}, while the analysis and derived physical properties are presented in \\S\\ref{sec:analysis}. We summarize our results in \\S\\ref{sec:conclusions}. Throughout the paper we assume $H_0=70$ \\kms, $\\Omega_M=0.27$, and $\\Omega_\\Lambda=0.73$, consistent with the five-year WMAP cosmological analysis \\citep{komatsu2009}. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we present \\irs\\ spectroscopic confirmation of a luminous infrared galaxy at $z=2.79$ that is highly magnified by the Bullet Cluster. We also spatially resolve the galaxy with \\wfc, revealing gravitational arcs near the previously determined \\spitzer\\ positions.  The redshift determination and high signal-to-noise of the spectrum, coupled with the multiwavelength photometry, enable us to refine our previous analysis of the physical properties of this galaxy. We determine that the lensed object is a low-mass dwarf galaxy, M$_*=4\\times10^{9}$ \\invmab M$_\\sun$, for which the far-infrared thermal emission is star-formation dominated. The inferred specific star formation rate, $SSFR\\approx 25-40$ Gyr$^{-1}$, is typical for a field galaxy of this mass at this redshift \\citep[e.g.][]{reddy2006}. The observed and derived physical quantities are summarized in Tables \\ref{tab:observed} and \\ref{tab:physical}.  By virtue of the extreme lensing -- the two brighest images have a combined magnification of roughly 100 -- this object thus provides a powerful laboratory in which to probe the conditions in low-mass, star-forming galaxies during an epoch when they are otherwise inaccessible. Indeed, this galaxy is one of the two lowest mass systems with \\irs\\ spectroscopy at $z>2$.\\footnote{The lensed LBG MS1512-cB58 has a factor of $\\sim2-4$ lower stellar mass \\citep[see][]{siana2008}.} Perhaps most promising for future investigations, our detection of rotational H$_2$ lines is indicative of a large molecular gas reservoir. We derive a temperature of $T=377^{+68}_{-85}$ K and estimate a gas mass of $M_{H_2}=2.2^{+17}_{-0.8}\\times10^8$~M$_\\sun$ in this component, roughly 2-42\\% of the stellar mass.  Future spectroscopic observations of longer wavelength H$_2$ lines with \\herschel\\ and CO rotational lines with a facility such as the Large Millimeter Telescope or the Atacama Pathfinder EXperiment (APEX) therefore have the potential to precisely measure the total molecular gas content in this galaxy.  Finally, we note that because dwarf galaxies of this mass are a ubiquitous population, the odds are good that programs like the \\herschel\\ Lensing Survey can detect a sizable sample of similar lensed sources behind massive clusters."
    },
    "1005/1005.4430_arXiv.txt": {
        "abstract": "Measuring the full  three-dimensional motions of extra-galactic objects in the Universe presents a seemingly insurmountable challenge. In this paper we investigate the application of a technique to measure tangential motion that has previously only been applied nearby within the Local Group of galaxies, to clusters of galaxies far beyond its borders. We show that mapping the mean line-of-sight motion throughout a galaxy cluster could in principle be used to detect the {\\it perspective rotation} induced by the projection of the cluster's tangential motion into the line-of-sight. The signal will be most prominent for clusters of the largest angular extent, most symmetric intrinsic velocity distribution and surveyed with the largest number of pointings possible. We investigate the feasibility of detecting this signal using three different approaches: measuring line-of-sight motions of individual cluster members; taking spectra of intracluster gas; and mapping distortions of the Cosmic Microwave Background radiation. We conclude that future spectroscopic surveys of 1000's of members of nearby galaxy clusters hold the most promise of measuring cluster tangential motions using this technique. ",
        "introduction": "It is our unhappy luck that even though we live in three dimensions, we are forced to view most of the Universe as a two dimensional projection with only one dimension of velocity information.  As such, we know very little about the motions of the things around us in any direction other than radially outward.  A more comprehensive view of motions in the Universe would provide a unique probe into the history and future of the things around us. Indeed, within the Milky Way, the HIPPARCOS \\citep{perryman97} and near-future GAIA \\citep{perryman02}  missions have brought about a rebirth in the ancient field of astrometry, pushing accuracies for direct proper motion measurements down first to mas/year precision and then on to tens of $\\mu$as/year. The prospect of billions of stars with full 6-D phase-spa<ce motion known for an appreciable fraction of the Galaxy has allowed astronomers to propose measuring Galactic structure \\citep[e.g.][]{johnston99} and recovering Galactic history \\citep[e.g.][]{helmi00} in unprecedented detail.  Similar measurements of the full  phase-space positions of objects throughout and beyond the Local Group would allow an analogous reconstruction of the masses and past interaction of these objects and test the existence of large-scale flows in the Universe \\citep[see][for an investigation on Local Group scales]{shaya03}. Cosmological simulations of structure formation suggest that the peculiar velocity distribution of galaxy clusters has root-mean-square spread of $\\sigma_{\\rm pec} \\sim$ 500 km/s \\citep{sheth01}, corresponding to proper motion scales for galaxy clusters at distance $d$ from us of \\begin{equation} \\mu \\sim \\left({v_{\\rm tan} \\over 500 {\\rm km/s}}\\right) \\left({10 {\\rm Mpc} \\over d}\\right) 10 {\\mu {\\rm as/year}}. \\label{eq:pm} \\end{equation} At first sight, this suggests that the next generation of astrometric missions would have sufficient precision to detect these transverse motions. However, prospects for direct measurements remain dim both because individual galaxies are extended and galaxy clusters have large internal velocity dispersions.  As an alternative to direct measurements of proper motions of nearby {\\it stellar} clusters, Galactic astronomers have traditionally employed a neat geometrical trick that allows them to infer the motion of an object in three-dimensions from line-of-sight velocities alone. The method relies on the fact that the projection of the transverse motion into the line-of-sight will induce a gradient in the average line-of-sight velocities measured across any extended object --- an effect known as known ``perspective rotation'' \\citep[hereafter PR ---][]{1961Feast}. PR has been used to verify measurements of the proper motion of the globular cluster Omega Centauri \\citep{1997Merritt}, as well as the distance to the Large Magellanic Cloud \\citep{2000Gould}. Most recently, \\citet{2008Kaplinghat} pointed out that there are now sufficient spectra taken for stars in nearby dwarf spheroidal galaxies for a precision of $\\sim$ 100 km/s in tangential velocity estimates to be possible. \\citet{walker08} subsequently verified this assertion with estimates for the motion of Fornax and Carina that agreed with prior astrometric measurments and the first three-dimensional measurement of the motion of Sextans. They also made estimates for Sculptor's proper motion that disagreed with prior work --- this disagreement could be explained as contamination by Sculptor's intrinsic rotation.  To date, M31 is the most distant object that PR has been measured for \\citep[using the line-of-sight velocities of 17 of its satellite galaxies, see][]{2008VanDerMarel} --- and this study represents over an order-of-magnitude leap in the distance to which such a measurement had been attempted. However, there is no limit in principle to the distance of objects for which this technique could prove useful. Indeed, many nearby clusters of galaxies have angular sizes of the same order-of-magnitude as M31's satellite system, with measurements of line-of-sight velocities of hundreds of members already taken and of ever increasing numbers of objects feasible in the near future.  In this paper, we examine to what extent we can expect to map bulk motions in the nearby Universe using current or near-future capabilities. We review the technique of PR in more detail in \\S 2, discuss 3 different types of data that should contain PR signatures in \\S 3 and summarize our conclusions in \\S 4. ",
        "conclusions": "We have outlined three possible ways to measure the PR of galaxy clusters and estimate their transverse motions: from line-of-sight velocity measurements of galaxy cluster galaxies, the motion of cluster gas and mapping the kSZ distortions in the CMB. All three of these measurements are unfeasible with the current data but could become feasible with near future instruments or larger data sets. The amplitude of PR is most significant for galaxy clusters with larger angular extent. Hence, the most promising approach with current capabilities is to survey as large a sample as possible of spectra of galaxies in nearby galaxy clusters.  With all approaches, the signal of PR could be confused by random (or turbulent) motions and intrinsic rotation of their system of targets. We are optimistic that it will be possible to disentangle these effects with sufficiently large data sets since PR imposes a unique signature of solid-body rotation on top of these sources of confusion. Indeed, \\citet{2008Kaplinghat} found that, when they included a small intrinsic rotation in their models of dwarf spheroidal galaxies, the error bars on the transverse velocity estimates derived by ``observing'' line-of-sight velocities of stars in their models increased by only a factor of two. While galaxy clusters are unlike dwarfs spheroidal galaxies in that they are not expected to be as relaxed or spherically symmetric, the \\citet{2008Kaplinghat} study provides a first step towards a comprehensive modeling effort.  Note that several other approaches to detecting transverse motions of galaxy clusters have also been proposed. For example using polarization maps of the CMB \\citep{sunyaev80}, gravitational lensing of the CMB \\citep{birkinshaw83}, and weak and strong lensing of background galaxies \\citep{molnar03}. The strength of the signatures in these methods are not dependent on the angular extent of the galaxy cluster and hence have the advantage over PR of being applicable to distant as well as nearby clusters.  With multiple possible directions for detection, we conclude that measurements of the full space motions of galaxy clusters are on the horizon."
    },
    "1005/1005.3210_arXiv.txt": {
        "abstract": "We report on an analysis of RXTE data of the transient neutron star low-mass X-ray binary (NS-LMXB) XTE J1701-462, obtained during its 2006-2007 outburst. The X-ray properties of the source changed between those of various types of NS-LMXB subclasses. At high luminosities the source switched between two types of Z source behavior and at low luminosities we observed a transition from Z source to atoll source behavior. These transitions between subclasses primarily manifest themselves as changes in the shapes of the tracks in X-ray color-color and hardness-intensity diagrams, but they are accompanied by changes in the kHz quasi-periodic oscillations, broad-band variability, burst behavior, and/or X-ray spectra. We find that the low-energy X-ray flux is a good parameter to track the gradual evolution of the tracks in color-color and hardness-intensity diagrams, allowing us to resolve the evolution of the source in greater detail than before and relate the observed properties to other NS-LMXBs. We further find that during the transition from Z to atoll, characteristic behavior known as the atoll upper banana can equivalently be described as the final stage of a weakening Z source flaring branch, thereby blurring the line between the two subclasses. Our findings strongly suggest that the wide variety in behavior observed in NS-LXMBs with different luminosities can be linked through changes in a single variable parameter, namely the mass accretion rate, without the need for additional differences in the neutron star parameters or viewing angle. We briefly discuss the implications of our findings for the spectral changes observed in NS LMXBs and suggest that, contrary to what is often assumed, the position along the color-color tracks of Z sources is not determined by the instantaneous mass accretion rate. ",
        "introduction": "\\label{sec:intro}  Neutron star low-mass X-ray binaries (NS-LMXBs) are systems in which a low-magnetic-field neutron star accretes matter from a low-mass companion star through Roche-lobe overflow. They display a wide range of X-ray spectral and variability properties, both as a class and as individual sources \\citep[see][for a recent review]{va2006}. Based on their correlated X-ray spectral and rapid variability behavior, NS-LMXBs are commonly divided into two subclasses, the so-called  Z sources and the atoll sources \\citep{hava1989}, named after the shapes they trace out in X-ray color-color (CD) and hardness-intensity diagrams (HID).  The Z sources have luminosities close to the Eddington luminosity ($\\sim$0.5--1\\,\\ledd\\ or more), whereas atoll sources have luminosities in the range $\\sim$0.01--0.5\\,\\ledd.  An example of the variety in shapes that is observed in the CD/HID tracks of the Z and atoll subclasses is shown in Figure \\ref{fig:overview}. The Z source tracks typically show three branches, which from top to bottom are called the horizontal branch, the normal branch, and the flaring branch. Based on the presence and orientation of these branches, the Z sources can be further divided into the `Cyg-like' Z sources (Cyg X-2, GX 5--1, and GX 340+0) and the `Sco-like' Z sources (Sco X-1, GX 17+2, and GX 349+2), with the former showing `Z'-shaped tracks in the HID (Fig.\\ \\ref{fig:overview}a) and the latter more `$\\nu$'-shaped tracks (Fig.\\ \\ref{fig:overview}b,c; see also \\citealt{kuvaoo1994,hovawi2007} (hereafter H07)). The atoll sources can show `Z'-shaped tracks as well (Fig.\\,\\ref{fig:overview}g), although the individual branches are thought to have a different physical nature than the Z source branches \\citep{baol2002,vavame2003,revava2004,va2006}. The branches of the Z-shaped atoll tracks are called, from top to bottom (Fig.\\ \\ref{fig:overview}g), the extreme island state, the island state, and the lower and upper banana branches; these branches are also frequently referred to as the hard, transitional/intermediate, and soft states, respectively. The low-luminosity atolls are typically found in the (extreme) island state, whereas the high-luminosity ones tend to be found mainly on the (upper) banana branch (Fig.\\ \\ref{fig:overview}d).  \\begin{figure}[h] \\vspace{0.5cm} \\epsscale{0.75} \\plotone{f1_compress.pdf} \\caption{Example CDs (left) and HIDs (right) of the Z source NS-LMXBs GX 5--1, GX 17+2, and GX 349+2, and the atoll source NS-LMXBs GX 9+9, 4U 1735--44, 4U 1636--53, and 4U 1705--44. Each data point represents a 256s average. The diagrams are taken from a compilation by Fridriksson et al.\\ (2010, in prep.) and were created following the steps described and referred to in \\S\\ref{sec:data} and \\S\\ref{sec:results}. Z and atoll branches/states are indicated in the CD or HID of each source. They are: horizontal branch (HB), normal branch (NB), flaring branch (FB), upper-banana branch (UB), lower-banana branch (LB), island state (IS), and extreme island state (EIS). For the Z source GX 5--1 we have also labelled the upturn from the horizontal branch (HB-upturn), and the `dipping' flaring branch (dipping FB).} \\label{fig:overview} \\end{figure}   In addition to motion along atoll or Z branches, some sources also show changes in the shape and location of their tracks in the CD and HID, which is often referred to as secular motion. Examples of this can be seen in Figure \\ref{fig:overview}, with the horizontal branch of GX 17+2 and the upper-banana branches of GX 9+9 and 4U 1735--44 being traced out at varying intensity levels. Much stronger secular motion can be found in the Z source Cyg X-2 \\citep{kuvava1996,wivaku1997}, where it occasionally results in switches between Cyg-like and Sco-like Z source behavior \\citep[see Figure 1 in][]{murech2002}. Finally, we note that transitions between the atoll (extreme) island state and the banana branches, and vice versa, often show strong hysteresis in HIDs, resulting in jumps between those branches  \\citep[e.g.,][see also Fig.\\ \\ref{fig:overview}f,g]{maba2004,gldogi2007,behomo2007}. It is not clear if and how this behavior is related to the secular motion.  In addition to differences in X-ray luminositiy and CD/HID tracks, the NS-LMXBs subclasses also show differences in X-ray spectral and variability properties \\citep{va2006}, radio behavior \\citep{mife2006}, as well as in the rate at which they show type I (thermonuclear) X-ray bursts \\citep{gamuha2008}. This disparity in the properties between the subclasses was a long-standing problem in the understanding of the accretion processes in NS-LMXBs. While differences in the mass accretion rate (\\mdot) could explain some of the variety in observed properties (e.g., in X-ray luminosity and bursting behavior),  it was unclear if one could link the wide variety in CD/HID tracks, rapid variability, and broadband spectral properties through changes in \\mdot\\ alone. Additional parameters have been proposed to explain some of the observed differences, not only between Z and atoll subclasses, but also between the Cyg-like and Sco-like Z sources: magnetic field strength and viewing angle \\citep[e.g.,][]{pslami1995,hava1989,kuva1995}.      Testing for the necessity of such additional parameters is complicated by the fact that none of the classical Z sources has been observed to venture into the atoll luminosity range (and vice versa). Until recently only two sources were suspected to have made a transition between Z-like and atoll-like behavior: Cir X-1 and XTE J1806--246 (= 2S 1803--245). Cir X-1 was found to show Z-like properties at high luminosities \\citep{shbrle1999}  and atoll-like properties at low luminosities \\citep{oovaku1995}. XTE J1806--246 showed some  characteristics of Z sources near the peak of its outburst, which reached $\\sim$\\ledd\\  (other bright NS-LMXB transients typically only reach $\\sim$0.5\\,\\ledd), and atoll properties at lower luminosities \\citep{wiva1999b}.  While these two examples suggested that a single source can show atoll/Z source characteristics depending on luminosity, and lent some credence to the belief that the NS-LMXB subclasses are entirely the result of differences in \\mdot, the actual evolution from Z tracks to atoll tracks had never been observed.  Significant progress on the issue of NS-LMXB subclasses was gained following the outburst of \\sj, a transient NS-LMXB that was discovered in January 2006 \\citep{reli2006} and remained in outburst for $\\sim$19 months. Initial observations with the {\\it Rossi X-ray Timing Explorer (RXTE)} revealed that \\sj\\ was the first transient NS-LMXB to show all of the Z source characteristics in X-rays (H07). The radio properties of the source were consistent with those of the Z sources as well \\citep{fedaho2007}. Observations of type I X-ray bursts with radius expansion, near the end of the outburst, put the source at a distance of 8.8$\\pm$1.3 kpc and indicate that during the early phase of the outburst the source had super-Eddington luminosities \\citep{lialho2009}. After the early phase of the outburst the source made a transition from  Cyg-like to Sco-like Z source behavior as the source intensity dropped by a factor of 1.5 (H07), and observations near the end of the outburst, in July 2007, showed the source evolving from Sco-like Z source behavior to atoll source behavior \\citep{hobewi2007,howial2007}.  A full spectral analysis of the entire outburst  was performed by \\citet[hereafter LRH09]{lireho2009}, who confirm that at the lowest luminosities the spectral behavior of \\sj\\ is consistent with that observed for the two atoll source transients Aql X-1 and 4U 1608--52 \\citep{lireho2007}. \\sj\\  is the first source in which the evolution from Z to atoll source has been observed in detail, and its behavior strongly suggests that different levels of \\mdot\\ are indeed sufficient to explain the different NS-LMXB subclasses.      The outburst of \\sj\\ also allows for a study of the evolution of the X-ray properties along CD/HID tracks within the context of the different NS-LMXB subclasses. This may shed new light on the nature of the Z and atoll branches, whether some or all are unique to each class, and, most importantly, what drives the motion along them. As far as atoll sources are concerned, past observations of transient NS-LMXBs have suggested that the evolution along atoll source CD/HID tracks is mainly driven by variations in \\mdot\\ \\citep[see, e.g.,][]{va1994b,lireho2007}, increasing from the extreme island state to the upper banana branch. However, the aforementioned hysteresis in spectral state transitions \\citep{maco2003} and  parallel tracks observed in HIDs \\citep[see, e.g.,][]{mevafo2001,dimeva2003} hint at a more complex relation between \\mdot\\ and spectral state in atoll sources \\citep{li2009}, not unlike what is suggested by the spectral evolution seen in transient black-hole X-ray binaries \\citep{hobe2005}.    It should be noted that there is ambiguity in some of these discussions about the definition of \\mdot.  Some authors implicity define \\mdot\\ as the mass transfer rate from the companion to the compact object Roche lobe, while others take the mass flow rate through the inner disk or the accretion rate onto the neutron star surface as the more relevant definition. Others explicitly distinguish various \\mdot\\ components.  In the presence of the various proposed flows in the accretion/ejection process, such as disk and equatorial boundary layer flows, spherical inflows, magnetically dominated polar flows, possible disk winds and jets, clearly this is a complex issue to which we shall return in \\S\\ref{sec:discussion}.  Interpreting the evolution along the Z source tracks has been more difficult, owing to the fact that motion along the Z source tracks corresponds to relatively small (less than a factor of $\\sim$2) and non-monotonic intensity variations. For Z sources it was suggested that \\mdot\\ increases from the horizontal branch to the flaring branch \\citep{havaeb1990}, although other scenarios have been put forward as well: e.g., \\mdot\\ changing in the opposite direction \\citep{chhaba2006}, or \\mdot\\ not changing at all \\citep{hovajo2002}. Based on their spectral analysis of the Z tracks, LRH09 suggest that motion along the Z tracks may occur at a nearly constant \\mdot, with the branches being the result of different instabilities in the near-\\ledd\\ accretion flow.      More detailed analyses of the data presented in this work, dealing with different aspects of \\sj, can be found in a number of other papers: spectral analysis (LRH09), type-I X-ray bursts \\citep{lialho2009}, kHz QPOs \\citep{sameal2010}, broad-band variability (Aresu et al.\\ 2010, in prep.), and the rapid decay into quiescence \\citep{frhowi2010}. Here, we present an overview of the outburst of \\sj\\ that focuses on aspects of the evolution of \\sj\\ in which the change from Z source to atoll source is most clearly seen: the CD/HID tracks, high-frequency QPOs, and broad-band variability.  In particular, we find that selections of data groups based on low-energy count rate allow for a more detailed study of the evolution of the CD/HID tracks than in H07 and LRH09. This also permits a more precise comparison of the observed tracks for \\sj\\ with those found in the various NS-LMXB sub-classes. In \\S\\ref{sec:data} we summarize our data set and analysis techniques. Our results are presented in \\S\\ref{sec:results} and in \\S\\ref{sec:discussion} these results are interpreted and discussed within the framework of various scenarios for the role of \\mdot\\ in driving the evolution between NS-LMXB subclasses and along the CD/HID tracks of NS-LMXBs. ",
        "conclusions": "The outburst of \\sj\\ has revealed in great detail the evolution of a low-magnetic field NS-LMXB from near- or super-Eddington rates all the way down to quiescence. Most interestingly, \\sj\\ was observed to transform from a Cyg-like Z source to a Sco-like Z source at high luminosities, and from Sco-like Z to atoll source at low luminosties. A brief comparison with other NS-LMXBs suggests that the behavior observed in \\sj\\ is representative of the entire class of low-magnetic field NS-LMXBs.  We conclude that these transformations between NS-LMXB subclasses are the result of changes in a mass accretion rate component \\mdot\\ (which for most of the outburst is measured by the 2--2.9 keV count rate) and argue that differences in neutron star spin/mass/magnetic field or viewing angle are of  minor importance to the Z/atoll phenomenology. The evolution observed within the Z source phase of the outburst shows a surprising evolution of the Z source tracks, with the horizontal, normal, and flaring branches disappearing one-by-one as \\mdot\\ appears to decrease. This is at odds with scenarios in which the same \\mdot\\ component increases or decreases monotonically from the horizontal to the flaring branch. We suggest that the \\mdot\\ responsible for the subclasses does not vary significantly along the Z tracks, while it does along atoll tracks. Motion along Z tracks may instead be the result of instabilities, whose presence {\\it does} depend on \\mdot."
    },
    "1005/1005.1645.txt": {
        "abstract": "We present the first sub-arcsecond resolution two-dimensional stellar kinematics and X-ray observations of the prototypical starburst galaxy NGC253 which define the position and nature of the galactic nucleus. These observations comprise some of the best probes of the central 300 pc of NGC253, the nearest massive galaxy undergoing a powerful starburst, and will allow us to gain more insight into the nature of the centers of starburst galaxies. %We present VLT/SINFONI near-IR integral-field spectroscopic observations and %\\textit{Chandra} imaging and spectroscopy of the central $300 \\times 300$ pc %region of the nearby starburst galaxy NGC253. %The near-IR $K$-band data are used to construct a two-dimensional stellar velocity %field from the $^{12}$CO absorption bandheads. %The kinematic center of the galaxy has been for the first time located from direct %measurements of the stellar kinematics in the nuclear region. We get an estimate of the stellar kinematic center location corresponding to an area of $r\\sim 1.2^{\\prime\\prime}$ centered $\\sim 0.7^{\\prime\\prime}$ southwest from the radio core, and historically presumed nucleus, TH2. %The infrared photometric center, or IR peak, lies $\\sim 2.6 \\arcsec$ away from the %kinematic center, and hence it is highly unlikely that it corresponds to the %galactic center. %This strong kinematic evidence clearly rules out the near-IR peak as the galactic %nucleus, lying $\\sim2.6 \\arcsec$ from the kinematic center Newly processed \\textit{Chandra} data reveal a central point-like hard X-ray source (X-1) lying $\\sim 0.4 \\arcsec$ southwest from the kinematic center. Very accurate alignment between radio, infrared and X-ray sources in the nuclear region shows that TH2, the IR photometric center and X-1 are not associated with each other. %being separated by $\\sim 1.1 \\arcsec$. As the kinematic center is consistent with the positions of TH2 and X-1, and both could be a manifestation of nuclear activity, we consider the two as possible galactic nucleus candidates. Although TH2 is the strongest compact radio source in the nuclear region, %brightest radio source and has the highest brightness temperature in the central region of NGC253. However, it does not have any infrared, optical or X-ray counterparts. If the kinematic center is associated with this source, by analogy we suggest that the nucleus of NGC253 resembles our Galactic Center SgrA$^*$. On the other hand, X-1 is a heavily absorbed object ($N_{H} = 7.5 \\times 10^{23}$ cm$^{-2}$) only detected at energies $>2$ keV ($L_{2-10 keV} \\sim 10^{40}$ ergs s$^{-1}$). %It is likely the X-ray counterpart of another strong radio and infrared source in %the central region. If X-1 is instead associated with the kinematic center, the nucleus of NGC253 is compatible with an obscured low luminosity active galactic nucleus (AGN) or a spatially resolved super star cluster (SSC) brightening up in X-rays most probably due to young supernovae or supernova remnants, a situation also observed in the nuclear starburst of M82. %we then suggest that the nucleus of NGC253 may harbor an obscured low luminosity %active galactic nucleus (AGN). %If X-1 corresponds to the kinematic center of the galaxy, %the simplest explanation accounting for these properties is that the nucleus of %NGC253 harbors a hidden LLAGN. %\\item Alternatively, X-1 may be associated with one SSC in this region: SSC %IR-11/radio-TH6. In this case the nucleus of NGC253 would then be a $10^5-10^6$ %M$_\\odot$ spatially resolved (diameter of $\\sim5$ pc) NSC in which some of its %massive stars are evolving into the SN phase. If no SSC is associated with the kinematic center, we conclude that NGC253 is a galaxy in which a strong starburst and a weak AGN (either TH2 or X-1) coexist. Results from few other high resolution studies of nearby starburst galaxies (e.g. M82, NGC5253, NGC4945) indicate that the AGN in these systems, if present, is always in the low luminosity regime. This may indicate that the onset of nuclear activity in galaxies is closely related with the occurrence of star formation, and that we are witnessing the emergence or dissaperance of an AGN. %Currently the gas in the nuclear region of NGC253 is feeding intense multiple %starbursts and the galactic nucleus is thus in a latent phase of activity.  %Both, the radio core and the central peak of hard X-ray emission, are consistent %with the kinematic center location, and therefore are the most likely galactic %three Ultraluminous X-ray sources (ULXs) in the central $10\\arcsec \\times %10\\arcsec$ but no source was detected at the position of TH2. %The strongest point-like X-ray source, HX-1, lies $\\sim 1.2 \\arcsec$ from TH2, has %an unabsorbed 2-10 keV luminosity of $\\sim 10^{43}$ ergs s$^{-1}$ and is observed %through an absorbing column density of $7.5 \\times 10^{23}$ cm$^{-2}$. %can be associated with another compact radio source TH6. %The two flat spectrum radio sources %We suggest that while HX-1 is either an intermediate-mass black hole or a weakly %accreting supermassive black hole, TH2 is a near twin of our Galactic Center %SgrA$^*$. ",
        "introduction": "\\label{int}  %In this paper we present the first sub-arcsecond resolution two-dimensional %stellar kinematics and X-ray observations of the prototypical starburst galaxy %NGC253 which define the position and nature of the galactic nucleus. %imaging observations of the 2.3 km CO absorption feature  It is observationally well established that star formation in starburst galaxies manifests itself in the production of numerous young massive star clusters of extraordinary luminosity and compactness, normally referred to as super star clusters (SSCs). In fact, the four most luminous circumnuclear starburst systems (M82, M83, NGC253 and NGC4945) account for $25\\%$ of the high-mass ($>8$ M$_\\odot$) star formation within 10 Mpc \\citep{heckman98}. Nearby starburst galaxies provide thus a unique laboratory for investigating the mechanisms which cause enhanced star formation to occur and be sustained, as well as physical conditions analogous to intensely star-forming galaxies identified at high-z (e.g. Shapley et al. 2003). An intriguing issue in studies of starburst galaxies is the apparent absence of a well defined nuclear source which could be associated with the center of the galaxy. %The key problem is that there exist offset between photometric centers (e.g. the %radio core is not coincident with the IR/visible photometric center). %A key problem in studies of starburst galaxies is to establish the physical nature %of the center of the galaxy. This situation is present in spatially resolved starbursts in nearby galaxies: e.g. M82 \\citep{matsumoto01}; NGC4945 \\citep{marconi00}; NGC5253 \\citep{summers04}; NGC253 (this work). Massive black holes (BHs) and/or nuclear star clusters (NSCs) are thought to reside in the centers of most galaxies, with the first being mostly identified in massive ellipticals \\citep{richstone98, graham08} and NSCs in irregular and late spirals (e.g. B\\\"oker et al. 2002; Seth et al. 2010). Although tens of SSCs exist in the central region of starburst galaxies, the presence of a NSC in these systems is very poorly constrained. Since the central few hundred parsecs of starburst galaxies are heavily affected by large and nonuniform dust extinction, the determination of SSCs properties (age, colors, luminosities, etc.), and thus the detection of a putative NSC, has been difficult and uncertain. %The ultimate observational evidence of the presence of massive BHs in the centers %of galaxies comes from kinematic studies, by measuring the gravitational influence %of the central object on neighboring stars and gas. As kinematic BH mass %determinations require very high spatial resolution, this has been possible to %achieve only in a few number of objects, being the most believable example the %supermassive black hole (SMBH) at the center of the Milky Way \\citep{begelman03}. The presence of massive BHs in galaxy centers is normally inferred when nuclear activity is observed \\citep{filippenko89}. %However, the presence or absence of any of them in starburst galaxies is very %poorly constrained. This is due to mainly two reasons: 1) the central few hundred %parsecs of starburst galaxies are heavily affected by large and nonuniform dust %extinction, making difficult and uncertain the determination of SSCs properties %and thus the detection of a putative NSC, and 2) the photometric centers at radio, %IR and X-ray wavelengths are not spatially coincident. The presence of massive BHs %in galaxy centers is normally inferred when nuclear activity is observed (ref %Seth). Active Galactic Nuclei (AGN) are usually characterized by prominent point sources in $K$-band. Starburst galaxies do not show any IR point-like emission at the position of the assumed galactic nucleus, which in the majority of these systems has been associated to either a powerful compact radio source or a hard X-ray peak. However, in several cases it is not possible to reject/accept the AGN hypothesis on the basis of radio/X-ray emission peaks alone, particularly when the three photometric centers (radio, IR, X-ray) don't coincide. For the nearest, best studied starbursts, possible galactic center candidates include: young supernovae (SNe), supernova remnants (SNRs), intermediate-mass BHs, inactive supermassive BHs, low-luminosity AGN, etc. (see e.g. Wills et al. 1999 and Stevens et al. 1999 for the case of M82, Marconi et al. 2000 for NGC4945 and Fern\\'andez-Ontiveros et al. 2009 for NGC253). %In addition, due to the strong, patchy obscuration in these systems the nature of %the SSCs and the photometric centers is far  As the center of a galaxy is likely to coincide with the bottom of the galaxy potential well, independent estimates of the kinematic center position from kinematic data can help to determine which photometric center, if any, is the most promising candidate for galactic nucleus. These important dynamical studies had not previously been widely performed because high S/N spectra over a full 2D field with high spatial resolution are required, and these are still scarce. In this paper we provide the first estimate of the kinematic center of the prototypical starburst galaxy NGC253 using sub-arcsecond resolution two-dimensional stellar kinematics. This is combined with data at radio, infrared and X-ray wavelengths to examine the nature of the galactic nucleus itself. These observations comprise some of the best probes of the inner region of NGC253 and will allow us to gain more insight into the nature of the centers of starburst galaxies. %and X-ray observations of the prototypical starburst galaxy %NGC253 which define %the position and nature of the galactic nucleus.  %In this paper we investigate the nature of the nucleus of the prototypical %starburst galaxy NGC253 using %existence of AGNs in starburst galaxies %through observations of NGC253, one of the closest and hence well studied %starburst galaxies, which is an ideal laboratory to test the power of integral %field spectroscopy and stellar kinematics in the determination of kinematic %centers of galaxies, particularly those which are largely inaccessible optically %because of dust obscuration, %has served as testing ground for study starburst theories %NGC253 is an ideal laboratory for studying any potential causal connection %between starburst and AGN activity. %provide an estimate of the kinematic center of the prototypical starburst galaxy  %NGC253. Understanding the nature of the nuclei of starburst galaxiesThis is a %crucial step in understanding how a putative massive BH and a strong starburst %relate to each other. %for studying any potential causal connection between starburst and AGN activity.  %With a strong patchy and with tens of SSCs detectable in the near-IR, the search %for a NSC in these systems is far from a simple task.   making difficult and %uncertain the detection of SSCs and thus the determination of their properties,  %Over the past decade the search for central massive objects (CMOs) in local %galaxies has experienced a major break-through. Mostly thanks to high spatial %resolution techniques and the development of new theoretical tools for the %analysis of stellar dynamical data, we have learned that probably all centers of %galaxies contain CMOs in the form of massive black holes (BHs) and/or nuclear star %clusters (see Richstone et al. 1998, Bender \\& Kormendy 2003, Seth et al. 2010). %, which in addition are largely inaccessible optically because of dust %obscuration,  %Over the past decade, there has been increasing interest in studying any possible %physical connection between circumnuclear starbursts and active galactic nuclei %(AGN). Thanks to high spatial resolution techniques it is now widely accepted that %many Seyfert galaxies contain both a circumnuclear starburst and an AGN (e.g. %Davies et al. 2007). To date, however, %little is known about the centers of nearby starburst galaxies ($<10$ Mpc) %it is not clear whether or not the centers of nearby starburst galaxies ($<10$ %Mpc) host massive black holes (BHs). %responsible for AGN activity. %or nuclear star clusters %or not starburst galaxies host accreting supermassive BHs.    %In the Milky Way, NGC4258 and likely M31, it is virtually certain that these dark %objects are indeed supermassive BHs, because alternatives in the form of stellar %remnants can be ruled out (Maoz 1998, Bender et al. 2005).  %The ultimate observational evidence of the presence of massive black holes (BH) in %the centers of galaxies comes from kinematic studies, by measuring the %gravitational influence of the central object on neighboring stars and gas. As %kinematic BH mass determinations require very high spatial resolution, this has %been possible to achieve only in a few number of objects, being the most %believable example the supermassive black hole at the center of the Milky Way %\\citep{begelman03}. %However, it is possible to perform kinematic studies at larger scales and estimate %the kinematic centers of galaxies from rotation curves, or even better, from %two-dimensional velocity fields. By associating the kinematic center with the %galactic nucleus, as it occurs in the majority of galaxies, one can then use other %methods to unveil the nature of the central object. Strong evidence of the %presence of a nuclear BH comes from its accretion-powered emission %\\citep{salpeter64}. Nuclear activity can be identified either from the detection %of broad recombination lines in the optical/near-IR spectrum of the nuclear %region, or from the presence of forbidden high ionization lines (coronal lines), %or from a point-like X-ray nuclear source, or a flat-spectrum radio core.     %As consequence, a key problem in studies of nearby starburst galaxies is to %establish the physical nature of the center of the galaxy. %, which in addition are largely inaccessible optically because of dust %obscuration, is to establish the physical nature of the center of the galaxy. %The fundamental question arising here is: do the centers of starburst galaxies %contain massive dark objects %, particularly when the photometric centers at radio, IR and/or X-ray wavelengths %don't spatially coincide  %Previous studies on the nature of the nuclei of starburst galaxies have not %reached a general consensus; cases were made for supernova remnants (NGC1569, %Westmoquette), SSCs, intermediate-mass black holes (NGC253, Weaver et al. 2002), %inactive supermassive black holes (M82, Wills et al. 1999; NGC253, %Fernandez-Ontiveros et al. 2009), low-luminosity AGN (M82, japones; NGC253, Turner %\\& Ho) , and powerful AGN (NGC4945, Marconi et al. 2000). %In particular, it is important to know if it is possible to hide an AGN, %contributing significantly to the bolometric emission, when optical to mid-IR %spectroscopy and imaging reveal only a starburst component.       %There are several ways in which one can determine the position of the center of a %galaxy. As nuclear activity in galaxies is likely to coincide with the bottom of %the galaxy potential well, a central compact radio source is a good indicator of %the center of a galaxy \\citep{hummel84}. %One can also determine the center by fitting ellipses to the surface brightness %distribution. For those galaxies which also show a well-defined nuclear source at %infrared or visible wavelengths, one can derive the central position by fitting a %Gaussian to the central source. Similarly, nuclear star clusters can be used to %locate the center (Matthews \\& Gallagher 2002). Each of these estimates yields a %photometric center. Another very important independent estimate can be obtained %from kinematic data by constructing a model for a rotation curve or velocity %field. %a tilted-ring model for a rotation curve or velocity field. %In this case a kinematic center is derived. Ideally, these different %determinations should all coincide. Indeed, for most galaxies the kinematic center %agrees %with the position of the radio continuum source and/or the photometric center as %derived from IR/visible imagery. However, in some cases, there exist offset %between photometric centers (the radio core is not coincident with the IR/visible %photometric center) or between the kinematic and photometric centers, this being %an indication of disturbances such as merging processes, spiral arms or strong %bars. When photometric centers do not coincide, as in the protoypical starburst %galaxy NGC253 \\citep{fernandez09}, the kinematic center position can help to %determine which photometric center, if any, is the most promising candidate for %galactic nucleus. In this paper we provide such independent estimate of NGC253's %center. %In the latter case, the discrepancies indicate disturbances such as merging %processes, spiral arms or strong bars.  %The ultimate observational evidence of the presence of massive black holes (BH) in %the centers of galaxies comes from kinematic studies, by measuring the %gravitational influence of the central object on neighboring stars and gas. As %kinematic BH mass determinations require very high spatial resolution, this has %been possible to achieve only in a few number of objects, being the most %believable example the supermassive black hole at the center of the Milky Way %\\citep{begelman03}.  %By associating the kinematic center with the %galactic nucleus, as it occurs in %the majority of galaxies, one can then use other %methods to unveil the nature of the central object. Strong evidence of the %presence of a nuclear BH comes from its accretion-powered emission %\\citep{salpeter64}. Nuclear activity can be identified either from the detection %of broad recombination lines in the optical/near-IR spectrum of the nuclear %region, or from the presence of forbidden high ionization lines (coronal lines), %or from a point-like X-ray nuclear source, or a flat-spectrum radio core.  As the nearest massive galaxy experiencing a strong starburst (3.94 Mpc, $1^{\\prime\\prime}$ = 19 pc, Karachentsev et al. 2003), %NGC253 is an ideal laboratory for studying any potential causal connection %between starburst and AGN activity. %Despite its proximity %(3.94 Mpc, $1^{\\prime\\prime}$ = 19 pc, Karachentsev et al. 2003), many aspects %of its nuclear region remain poorly understood. One of them is precisely the %exact location and physical nature of the center of the galaxy. NGC253 is one of the most studied objects in the sky, yet many aspects of its nuclear region remain poorly understood. %despite its proximity. One of the main unresolved issues is the precise location and physical nature of the center of the galaxy. %A variety of investigations of NGC 253 have provided information about the %individual compact radio sources %located within the inner 100 pc of the galaxy \\citep{turner85, ulvestad97, %lenc06, brunthaler09}. The galactic nucleus has historically been associated to a powerful compact non-thermal radio source (TH2 in Ulvestad \\& Antonucci 1997) with characteristics similar to other synchrotron sources in active galaxies \\citep{turner85}: very small size ($<100$ mas), %(VLA A-configuration observations do not resolve TH2 at any wavelength, see %Table 3), %e.g. size $<130 \\times 70 mas$ at 1.3 cm), high brightness temperature ($T>40000$ K at 15 GHz), and a relatively flat spectral index between $8.4-23$ GHz($\\alpha =-0.3$, $S_\\nu \\propto \\nu^\\alpha$). Follow up studies accross the electromagnetic spectrum, each limited by the spatial resolution available at that time, have matched this source with i) H$_2$O maser emission \\citep{nakai95}, ii) different knots of IR emission \\citep{forbes00, galliano05}, and iii) the hard X-ray emission peak \\citep{weaver02}, %and the dynamical center of the galaxy (Arnaboldi, Yo), suggesting that TH2 may be a low luminosity active galactic nucleus (LLAGN). However, it is now apparent from recent high-spatial resolution observations at different wavelengths, and using diverse observational techniques, that this picture may be incorrect. The evidences are the following:  \\begin{itemize} \\item VLA and VLBA observations at 23 GHz \\citep{brunthaler09} have shown that the high velocity H$_2$O maser is not associated with TH2 but with another nearby continuum source (TH4 in Ulvestad \\& Antonucci 1997), which has been classified as a SNR embedded in an HII region \\citep{lenc06}. This suggests that a maser-AGN connection is unlikely and that the maser emission is almost surely related to star formation.  \\item Very high spatial resolution VLBI observations at 2.3 and 23 GHz \\citep{lenc06, brunthaler09} did not detect TH2 on mas scales. As this source is unresolved and very bright in VLA 5-23 GHz continuum maps (beam sizes $>100$ mas), %this implies that the source is either over-resolved or is too faint. In both %cases this non-detection questions this questions the presence of a possible AGN. %However, as pointed out by \\citet{brunthaler09}, this non-detection does not rule %out other scenarios such as the existence of an inactive supermassive black hole %like Sgr A$^*$.  \\item Very accurate alignment at parsec scales has been found between the near-IR and radio sources by means of Adaptive Optics VLT/NACO IR data with a comparable resolution to those in existing VLA radio maps \\citep{fernandez09}. Surprisingly, this new alignment shows that TH2 has no infrared counterpart. The newly revealed IR nature of TH2 shows no evidence of a central point source that would indicate nuclear activity. However, the non-detection of IR emission does not rule out other scenarios such as the existence of an inactive supermassive black hole (SMBH) like Sgr A$^*$ \\citep{fernandez09, brunthaler09}. %and instead suggests that TH2 may be a twin of Sgr\\,A$^*$ rather than an AGN.  \\end{itemize}  In addition, one needs to be cautious when associating the hard X-ray peak to TH2. It is well known that NGC253 harbors $\\sim 60$ compact radio sources \\citep{ulvestad97} and $\\sim37$ near-infrared point-like sources \\citep{fernandez09}. This significant change in morphology with wavelength has led to different registrations in the literature (see e.g. Galliano et al. 2005 and references herein). Since \\textit{Chandra} data were processed with $0.5\\arcsec$ pixel $^{-1}$ resulting in $\\sim 1 \\arcsec$ resolution and have an astrometric accuracy of $\\sim 1 \\arcsec$ rms, the hard X-ray peak could be associated with several other discrete sources in the field. %or simply just not coincident with TH2. This is demonstrated in Figure 2 of \\citet{weaver02} where a variety of IR and radio sources are contained within the size of the hard X-ray peak. Moreover, these authors followed the alignment proposed by \\citet{forbes00} which, in light of the current high spatial resolution NACO images presented in \\citet{fernandez09}, seems to be incorrect. Thus, it is still a topic of active investigation whether TH2 harbors an AGN, or even if it corresponds to the galactic center, or not.  Based on the alignment proposed by \\citet{fernandez09} one could point to the infrared photometric center (or IR peak) as possible AGN candidate. The IR peak lies $\\sim 3.3 \\arcsec$ SW from TH2. It is in fact the brightest source at infrared wavelengths ($1-20\\, \\mu$m) and emits $\\sim10\\%$ of the galaxy's bolometric flux \\citep{mohan02}. The spatial distribution of the 3.3 $\\mu$m polycyclic aromatic hydrocarbons (PAH) feature-to-continuum ratio reaches a local minimum in this region \\citep{tacconi05}. This indicates that the source produces strong ionization, probably destroying the PAH molecules by photo-dissociation. Supporting this interpretation is the fact that multiband tracers of ionized gas such as Pa$\\beta$, Br$\\gamma$, [HeI] and [NeII] are seen to peak strongly at this position \\citep{boeker98, engelbracht98, keto99, forbes00, tacconi06}. %which is in contrast with the behaviour of other nuclear sources. In addition, the IR peak has an individual counterpart at radio wavelengths (TH7 in Ulvestad \\& Antonucci 1997) which also presents a relatively flat spectrum between $8.4-23$ GHz ($\\alpha=-0.21$, $S_\\nu \\propto \\nu^\\alpha$), but it is $\\sim 4$ times fainter than TH2. %It has been interpreted in terms of a mixture distribution of supernova remnants %and HII regions \\citep{ulvestad97}. Moreover, this strong IR source is located at the vertex of the rather collimated ionised-gas region seen in NGC253 (see Figure 8 of Tacconi-Garman et al. 2005). All these characteristics make this object unique among the other sources in the central region of the galaxy and could be explained either by invoking a very young SSC or an AGN. %The IR peak has been thought to host a Super Star Cluster (SSC). %If there is an AGN in this galaxy, these characteristics point directly to this %source as a candidate.  NGC253 has been the subject of many detailed kinematical studies of its gaseous components. Optical rotation curves of the ionized gas have shown that, in addition to ordered rotation, large non-circular motions exist within the central region \\citep{ulrich78, munoz93, arnaboldi95}. However, because of problems of obscuration at optical wavelengths, results from H$\\alpha$ and [N II] observations are difficult to interpret. In addition, HI and H$\\alpha$ rotation curves are incomplete: either they cover the outer parts $r>100\\arcsec$ \\citep{pence80}, or they are confined to the innermost regions where peculiar features and asymmetries in the ionized gas are dominant. Most of the optical and IR observations, aimed at studying the kinematics of the molecular gas in the nuclear region of NGC 253 (e.g., Prada et al. 1996; Engelbracht et al. 1998), have relied on measuring the velocity gradients along chosen position angles passing through the nucleus (e.g., major axis and minor axis). Interferometric radio observations at millimeter and centimeter wavelengths measure two-dimensional velocity fields which provide a more complete picture of the kinematics in the nuclear region. Several radio observations in molecular emission lines such as CO \\citep{mauersberger96, canzian88, das01, paglione04}, CS \\citep{peng96}, HCN \\citep{paglione95}, or in hydrogen recombination lines \\citep{anantamariah96} exist in the literature. These two-dimensional measurements reveal a complex but systematic velocity pattern in the central region. In addition to solid-body rotation, the observed velocity fields indicate motions which may be due to a barlike potential in the center \\citep{peng96} or a kinematic subsystem which may be caused by a past merger event (Anantharamaiah \\& Goss 1996). Although most of these observations have high spectral resolution, they typically have poor spatial resolution $>2.5 \\arcsec$ and are not always useful for probing the dynamics very close to the nucleus.  Because of this situation, different authors assumed different centers for their rotation curves/velocity maps: either the center is derived assuming complete symmetry of the outer parts in their fields of view, or it is coincident with the IR peak, or with the radio core TH2, or with the surface brightness peak. The highest resolution measurement of the two-dimensional velocity field available to date is that of Anantharamaiah \\& Goss (1996), who observed the H92$\\alpha$ recombination line with a beam of $1.8\\arcsec \\times 1.0 \\arcsec$ and a velocity resolution of 54 km s$^{-1}$. After trying several positions near the H92$\\alpha$ emission peak, the optimum center of the velocity field was found to be at ($\\alpha$, $\\delta$)$_{2000}$ = $00^h47^m33^s.06$, $-25\\degr 17^\\prime 18\\arcsec .3$ ($1\\sigma$ error $\\sim 0.3\\arcsec$), which is $\\sim 2.0 \\arcsec$ southwest of the H92$\\alpha$ line peak which in turn is coincident with TH2. This means that the most accurate estimate of the kinematic center based on gas kinematics available to date lies closer to the IR peak ($\\sim 1.3\\arcsec$) than to TH2. %is not consistent with TH2, but lies closer to the IR peak within their %resolution. %could be associated within the errors with the IR peak.  However, kinematic studies of galaxies focused on the gaseous component (ionized and/or molecular gas) have the potential for revealing the dynamical forces that work in the gas closer to the galactic nucleus, but they cannot trace adequately the gravitational potential. The gas represents only a fraction of the total mass of the galaxy and could be driven by non-gravitational perturbations such as interactions, warps, bars, etc. Therefore, the kinematic parameters of any nuclear galactic disk (center, position angle and inclination) can be greatly influenced by these perturbations. Thus the velocity field of stars is needed to properly trace the gravitational potential.  Integral-field spectroscopy is an ideal tool for studying objects with complex morphologies and kinematics, such as the nuclei of starburst galaxies. This paper is the first of two describing the conclusions from an integral-field spectroscopic study with SINFONI on the VLT to clarify the structure, kinematics and emission mechanisms of the central 150 pc starburst region of NGC253. In this part we use the CO stellar absorption features in $K-$band to determine the stellar kinematic center of the galaxy. The stellar CO features have the advantage that they trace a component of the galaxy nucleus that is not affected by non-gravitational processes such as shocks or winds. %In order to reveal its nature, We also obtained archival \\textit{Chandra} data and re-processed them using $0.125\\arcsec$ pixel$^{-1}$. The X-ray images are carefully compared with available high spatial resolution radio and near-IR data. This multiwavelength approach combined with the estimate of the kinematic center is essential to find the true nucleus of NGC253. %We carefully aligned this new X-ray images with high spatial resoultion radio %and near-IR images. As a result we will be able to identify individual %counterparts to the %As neither broad recombination lines nor forbidden high ionization lines have %been observed in the nuclear spectra of this galaxy, AGN activity can be %identified  from a point-like X-ray nuclear source, or a flat-spectrum radio %core. In a forthcoming publication \\citep{mueller09} we investigate the physical properties and kinematics of the super star clusters and the molecular and ionized gas in the nuclear starburst. %We address the star formation process, stellar populations, masses of individual %star clusters and the kinematics of the superwind.  This paper is structured as follows. In Section 2, we describe the SINFONI and \\textit{Chandra} observations and the processing we performed on the data. The method used to extract the kinematics from the SINFONI spectral features is presented in this Section as well as the new X-ray images and spectra. In Section 3 we examine the stellar velocity field and obtain an estimate of the position of the kinematic center. We discuss the implications of our results for the nature of the galactic nucleus in Section 4. The overall conclusions of the study are outlined in Section 5. ",
        "conclusions": "\\label{conclusions}  SINFONI near-infrared integral field spectroscopy combined with \\textit{Chandra} imaging and spectroscopy have been used in this work to probe the galactic nucleus of NGC253. Thanks to our high S/N integral field data we were able to extract 2D stellar kinematics in the nuclear region of NGC253 and estimate the kinematic center of the galaxy with a very good accuracy. Our main results and conclusions can be summarized as follows:  \\begin{itemize} \\item  The stellar kinematic center is located in between the strongest compact radio source (TH2) and the IR photometric center, or IR peak. It lies $\\sim 0.7 \\arcsec$ SW from TH2 and $\\sim 2.6 \\arcsec$ NE from the IR peak. \\item The estimated confidence interval of our kinematic modeling ($r\\sim 1.2\\arcsec$) clearly discards the IR peak as kinematic center, but includes the position of TH2. \\item Newly processed \\textsl{Chandra} images reveal a central point-like hard X-ray source (X-1) lying $\\sim0.4 \\arcsec$ SW from the kinematic center. %and two off-center X-ray sources which are good Ultraluminous X-ray sources (ULX) %candidates. \\item Very accurate alignment between radio, IR and X-ray sources in the nuclear region shows that TH2 and X-1 are not associated with each other, being separated by $\\sim 1.1 \\arcsec$. As the two could be a manifestation of nuclear activity, and our current estimate of the kinematic center location is consistent with both sources, we consider the two as possible galactic nucleus candidates. \\item TH2 is the brightest radio source, with the highest brightness temperature, detected with the VLA in the central region of NGC253. However, it does not have any IR, optical or X-ray counterparts. When observed at very high spatial resolution with the VLBI, TH2 is among the weakest sources ($\\sim 6$ mJy at 2.3 GHz, Lenc \\& Tingay 2006), or it is not detected ($<1$ mJy at 23 GHZ, Brunthaler et al. 2009). If the kinematic center is associated with TH2, we conclude that the nature of the nucleus of NGC253 resembles that of SgrA$^*$. \\item X-1 is a heavily absorbed object ($N_{H} > 10^{23}$ cm$^{-2}$) only detected at energies $>2$ keV ($L_{2-10 keV} \\sim 10^{40}$ ergs s$^{-1}$). If X-1 corresponds to the kinematic center of the galaxy, the simplest explanation accounting for these properties is that the nucleus of NGC253 harbors a hidden LLAGN. \\item Alternatively, X-1 may be associated with one SSC in this region: SSC IR-11/radio-TH6. In this case the kinematic center of NGC253 would then be a $10^5-10^6$ M$_\\odot$ spatially resolved (diameter of $\\sim5$ pc) NSC in which some of its massive stars are evolving into the SN phase. \\end{itemize}  %This study demonstrated the power of integral field spectroscopy and stellar %kinematics for the determination of kinematic centers of galaxies. %suggesting that this source is likely to be a SSC. %In any of the two cases, we can conclude that there is a massive central object %presenting signatures of low accretion-powered activity. %This allows us to conclude %that NGC253 is basically a galaxy in which a strong starburst %and a weak AGN (either TH2 or X-1) coexist. We speculate that the galactic nucleus %of NGC253 is in an evolutionary state where it is transitioning between a %starburst phase and an AGN phase. Currently the intense multiple starbursts are %young with ages $<35$ Myr \\citep{mueller09}. Previous studies of starburst-AGN %connection hint at a delay of $50-100$ Myr between the onset of star formation and %subsequent fuelling of the black hole \\citep{davies07}. The galactic nucleus of %NGC253 is thus in a latent phase of activity, still lacking of sufficient fuel. An important conclusion from this study is that the nature of dusty galaxies cannot be always characterized by photometric data alone but must be complemented with spectroscopic observations. In particular, we demonstrated the power of integral field spectroscopy for the determination of kinematic centers of galaxies. Further near-IR integral-field observations with high spatial resoultion would be useful to derive a more accurate position of the kinematic center, measure directly the mass of the putative central massive object, and reveal the presence of either broad recombination lines or forbidden high ionization lines. A detection of any of these would provide definite evidence of the existence of an AGN in this galaxy.  %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the"
    },
    "1005/1005.3951_arXiv.txt": {
        "abstract": "The emission of neutrinos within a wide energy range is predicted from very-high-energy phenomena in the Universe.  Even the current or next-generation Cherenkov neutrino telescopes might be too small to detect the faint fluxes expected for cosmic neutrinos with energies exceeding the EeV scale. The acoustic detection method is a promising option to enlarge the discovery potential in this highest-energy regime.  In a possible future deep-sea detector, the pressure waves produced in a neutrino interaction could be detected by a $\\gtrsim 100\\,$km$^3$-sized array of acoustic sensors, even if it is sparsely instrumented with about 100\\,sensors/km$^3$.  This article focuses on the AMADEUS set-up of acoustic sensors, which is an integral part of the ANTARES detector. The main aim of the project is a feasibility study towards a future acoustic neutrino detector. However, the experience gained with the ANTARES-AMADEUS hybrid opto-acoustic set-up can also be transferred to future very large volume optical neutrino telescopes, especially for the position calibration of the detector structures using acoustic sensors. ",
        "introduction": "At ultra-high energies (UHE, $E \\gtrsim 10^{18}\\,$eV), neutrinos are the only viable messengers beyond the ``local'' universe.  Photons interact with the microwave and infrared background light. Protons suffer energy losses due to photo-production of pions (GZK mechanism \\cite{grei,zats}) and of e$^+$e$^-$ pairs; nuclei are additionally subject to photo-nuclear reactions. All these reactions confine undisturbed propagation of UHE cosmic rays to distances below 100\\,Mpc \\cite{allard}.  UHE neutrinos, on the other hand, are undisturbed and ``guaranteed'' by the GZK mechanism. The interaction rate of neutrinos originating from the propagation of UHE protons in 1\\,km$^3$ water equivalent has been estimated to be about 0.2 per year \\cite{seckel}.  Consequently, target masses exceeding 100\\,km$^3$ are required to obtain a few neutrino events per year. Such detector sizes are not reachable with current technologies and new detection techniques have to be considered for the study of UHE cosmic neutrinos, e.g.\\ the acoustic method discussed here.  The investigation of acoustic neutrino detection has historically evolved in the context of Cherenkov neutrino {te\\-le\\-scopes}. The method was discussed in the 1970s within the DUMAND project \\cite{bowen}. Acoustic set-ups have been integrated in the framework of existing Cherenkov experiments in the 2000s: AMADEUS in {ANTA\\-RES}, ONDE in NEMO, SPATS in IceCube and acoustic test set-ups in the lake Baikal neutrino telescope\\footnote{There is a variety of additional ongoing projects - not only in sea water but also in ice and salt, cf.\\ \\cite{arena} for an overview.}. A next step towards an acoustic neutrino detector could be the integration of acoustic sensors into a next-generation neutrino telescope like KM3NeT \\cite{km3net}. ",
        "conclusions": "Cosmic neutrinos are the only viable UHE messengers beyond the ``local'' universe.  The acoustic neutrino detection technique is an option to extend neutrino astrophysics to the extreme energy range, either in stand-alone detectors or in combined opto-acoustic sensor arrays.  The AMADEUS system, dedicated to the investigation of this technique, has been successfully integrated into the ANTA\\-RES neutrino telescope.  Except for size, the system has all features required for an acoustic detector and thus allows for deciding on the feasibility of neutrino detection with a potential future large acoustic sensor array. In the context of AMADEUS it is currently being investigated whether an integrated opto-acoustic sensor configuration has the potential to combine the position calibration of an optical detector with acoustic detection studies and marine science applications."
    },
    "1005/1005.1679_arXiv.txt": {
        "abstract": "In this paper we test a special-relativistic formulation of Smoothed Particle Hydrodynamics (SPH) that has been derived from the Lagrangian of an ideal fluid. Apart from its symmetry in the particle indices, the new formulation differs from earlier approaches in its artificial viscosity and in the use of special-relativistic ``grad-h-terms''. In this paper we benchmark the scheme in a number of demanding test problems. Maybe not too surprising for such a Lagrangian scheme, it performs close to perfectly in pure advection tests. What is more, the method produces accurate results even in highly relativistic shock problems. ",
        "introduction": "\\label{Rosswog::sec:1} Relativity is a crucial ingredient in a variety of astrophysical phenomena. For example the jets that are expelled from the cores of active galaxies reach velocities tantalizingly close to the speed of light, and motion near a black hole is heavily influenced by space-time curvature effects. In the recent past, substantial progress has been made in the development of numerical tools to tackle relativistic gas dynamics problems, both on the special- and the general-relativistic side, for reviews see \\cite{Rosswog:marti03,Rosswog:font00,Rosswog:baumgarte03}. Most work on numerical relativistic gas dynamics has been performed in an Eulerian framework, a couple of Lagrangian smooth particle hydrodynamics (SPH) approaches do exist though.\\\\ In astrophysics, the SPH method has been very successful, mainly because of its excellent conservation properties, its natural flexibility and robustness. Moreover, its physically intuitive formulation has enabled the inclusion of various physical processes beyond gas dynamics so that many challenging multi-physics problems could be tackled. For recent reviews of the method we refer to the literature \\cite{Rosswog:monaghan05,Rosswog:rosswog09b}. Relativistic versions of the SPH method were first applied to special relativity and to gas flows evolving in a fixed background metric \\cite{Rosswog:kheyfets90,Rosswog:mann91,Rosswog:mann93,Rosswog:laguna93a,Rosswog:chow97,Rosswog:siegler00a}. More recently, SPH has also been used in combination with approximative schemes to dynamically  evolve space-time \\cite{Rosswog:ayal01,Rosswog:faber00,Rosswog:faber01,Rosswog:faber02b,Rosswog:oechslin02,Rosswog:faber04,Rosswog:faber06,Rosswog:bauswein10}.\\\\ In this paper we briefly summarize the main equations of a new, special-relativistic SPH formulation that has been derived from the Lagrangian of an ideal fluid. Since the details of the derivation have been outlined elsewhere, we focus here on a set of numerical benchmark tests that complement those shown in the original paper \\cite{Rosswog:rosswog09d}. Some of them are ``standard'' and often used to demonstrate or compare code performance, but most of them are more violent---and therefore more challenging---versions of widespread test problems. ",
        "conclusions": "We have summarized a new special-relativistic SPH formulation that is derived from the Lagrangian of an ideal fluid \\cite{Rosswog:rosswog09d}. As numerical variables it uses the canonical energy and momentum per baryon whose evolution equations follow stringently from the Euler-Lagrange equations. We have further applied the special-relativistic generalizations of the so-called ``grad-h-terms'' and a refined artificial viscosity scheme with time dependent parameters.\\\\ The main focus of this paper is the presentation of a set of challenging benchmark tests that complement those of the original paper \\cite{Rosswog:rosswog09d}. They show the excellent advection properties of the method, but also its ability to accurately handle even very strong relativistic shocks. In the extreme shock tube test 3, where the post-shock density shell is compressed into a width of only 0.1 \\% of the computational domain, we find the shock front to propagate at slightly too large a pace. This artifact ceases with increasing numerical resolution, but future improvements of this point would be desirable. We have further determined the convergence rate of the method in numerical experiments and find it first-order accurate when shocks are involved and second-order accurate for smooth flows.  \\vspace*{0.5cm}  {\\em Acknowledgment}\\\\ This work was supported by the German Research Foundation under grant number 50245 DFG-RO-5."
    },
    "1005/1005.5570_arXiv.txt": {
        "abstract": "We present the second report of our systematic search for strongly lensed quasars from the data of the Sloan Digital Sky Survey (SDSS). From extensive follow-up observations of 136 candidate objects, we find 36 lenses in the full sample of 77,429 spectroscopically confirmed quasars in the SDSS Data Release 5. We then define a complete sample of 19 lenses, including 11 from our previous search in the SDSS Data Release 3, from the sample of 36,287 quasars with $i<19.1$ in the redshift range $0.6<z<2.2$, where we require the lenses to have image separations of $1''<\\theta<20''$ and $i$-band magnitude differences between the two images smaller than $1.25$~mag. Among the 19 lensed quasars, 3 have quadruple-image configurations, while the remaining 16 show double images. This lens sample constrains the cosmological constant to be $\\Omega_\\Lambda=0.84^{+0.06}_{-0.08}({\\rm stat.})^{+0.09}_{-0.07}({\\rm syst.})$ assuming a flat universe, which is in good agreement with other cosmological observations. We also report the discoveries of 7 binary quasars with separations ranging from 1\\farcs1 to 16\\farcs6, which are identified in the course of our lens survey. This study concludes the construction of our statistical lens sample in the full SDSS-I data set. ",
        "introduction": "\\label{sec:intro}  Gravitationally lensed quasars are useful tools for a variety of astrophysical and cosmological studies \\citep[e.g.,][]{turner84,blandford87,schneider92,kochanek06}. In particular, statistical analyses of lensed quasars serve as a useful probe for the cosmological constant \\citep{fukugita90,turner90} and the Hubble constant \\citep{oguri07a}. The Hubble Space Telescope ({\\em HST}) Snapshot survey \\citep{maoz93} and the Cosmic-Lens All Sky Survey \\citep[CLASS;][]{myers03,browne03} have provided examples of complete lens samples that have allowed cosmological studies. The {\\em HST} Snapshot survey includes five lenses selected from 502 bright, relatively high-redshift quasars, and was used to derive a limit on the cosmological constant \\citep{maoz93}. It was also used to study galaxy evolution by combining with other lens surveys \\citep{chae10}. CLASS contains 13 lenses from 8958 radio sources; the redshift distribution of this sample is not well determined \\citep[e.g.,][]{munoz03}. This sample has also been used to constrain cosmological models as well as the structure and evolution of lens galaxies \\citep[e.g.,][]{rusin01,mitchell05,chae06}. While large samples of galaxy-galaxy lenses are being assembled by various groups \\citep[e.g.,][]{bolton06,bolton08,cabanac07,faure09,marshall09,kubo09,feron09} and are used to study the structure of the lens galaxies, they are not well suited as a cosmological probe because galaxy-galaxy lenses are often selected from samples of lensing objects, not source objects, much complicating the statistics.  In order to construct a large statistical sample of lensed quasars that can be used as a cosmological probe, we have conducted the Sloan Digital Sky Survey Quasar Lens Search \\citep[SQLS;][hereafter Paper I]{oguri06} based on a sample of spectroscopically confirmed quasars \\citep{schneider07} derived from the Sloan Digital Sky Survey \\citep[SDSS;][]{york00}.  To allow cosmological tests, the lens sample must be complete with well-defined criteria. We refer to a sample that allows statistical tests as a ``statistical sample''.  Our lens sample is, in fact, designed to be complete under prescribed conditions and therefore suitable for statistical studies and cosmological tests, given the accurately defined selection function (see Paper I) and the homogeneity of the SDSS data.  In \\citet{inada08}, hereafter Paper II, we presented a complete sample of 11 lensed quasars suitable for statistical analyses, selected from the 22,683 quasars satisfying $0.6<z<2.2$ and $i<19.1$\\footnote{Here $i$ is the Point Spread Function (PSF) magnitude corrected for Galactic extinction from the maps of \\citet{schlegel98}.}  out of the total of 46,420 quasars \\citep{schneider05} in the SDSS Data Release 3 \\citep[DR3;][]{abazajian05}. This sample derived from the DR3 gives cosmological constraints \\citep[][hereafter Paper III]{oguri08a} that agree with the current cosmological model \\citep[e.g.,][]{komatsu09,tegmark06}.  In this paper, we extend our source population to the SDSS Data Release 5 \\citep[DR5;][]{adelman07}, concluding the SDSS-I (the first phase of the SDSS project through 2005 June). The selection process is the same as that used to create the DR3 statistical sample of lenses, as described in Paper II. Note that all lensed quasars described in this paper have already been reported in their discovery papers. The main focus of this paper is to define the DR5 statistical sample, and to report quasars that turned out not to be gravitational lenses. We briefly describe the source quasar sample from which we constructed the DR5 statistical lens sample in Section~\\ref{sec:source}, and the selection of lens candidates in Section~\\ref{sec:cand}. We present observational results for the lensing candidates in Section~\\ref{sec:obs}, and the resulting DR5 lensed quasar sample and updates of cosmological constraints in Section~\\ref{sec:dr5_lens}. Section~\\ref{sec:summary} gives a summary of our results. ",
        "conclusions": "\\label{sec:summary}  We have completed our systematic lensed quasar search in the SDSS-I quasar sample presented in the DR5 quasar catalog \\citep{schneider07}. With follow-up observations for 136 lens candidates, we found ten lensed quasars beyond our DR3 sample. Eight of them, including one previously known, are catalogued in our complete statistical lensed quasar sample within specified separation ranges and magnitude differences. These conditions minimize our selection bias for lensed systems.  Combining with the result from DR3, we present a complete sample of 19 lensed quasars selected from 36,278 quasars, where three of them have quadruple-image configurations. We then updated cosmological constraints obtained in Paper III, and obtained results consistent with other cosmological observations \\citep[e.g.,][]{komatsu09}.  In addition to the 19 lensed quasars in our complete sample, the DR5 quasar catalog contains at least 17 additional lenses. Two were discovered among the 136 lens candidates considered here but excluded from our complete lens sample, and two were discovered from quasars outside the source sample of the 36,278 quasars. The remaining 13 lenses include two previously known lenses and 11 lenses from the DR3 sample (Paper II). These numbers may be compared with those of the CLASS, which contains 13 lenses in their statistical sample with well defined criteria and 22 lenses in total among 8958 quasars.  In this paper, we have also report the discovery of 7 binary quasars with nearly identical redshifts, as well as 8 projected quasar pairs. These quasar pairs are a useful addition to the studies of the small-scale correlation function and interaction of quasars \\citep{hennawi06a,myers08,green10} and the spatial distribution of absorbers \\citep{bowen06,hennawi06b,tytler09}.  The quasar sample used in the present work is constructed from the full SDSS-I data set, and hence the work represents the completion of our statistical lens sample in the SDSS-I. We plan to continue our lens survey further to construct the SQLS lens sample from the SDSS-II, using the DR7 quasar catalog \\citep{schneider10}. More detailed analysis of cosmological constraints will be presented elsewhere."
    },
    "1005/1005.0814_arXiv.txt": {
        "abstract": "{In this paper we study possible observational consequences of the bouncing cosmology. We consider a model where a phase of inflation is preceded by a cosmic bounce. While we consider in this paper only that the bounce is due to loop quantum gravity, most of the results presented here can be applied for different bouncing cosmologies. We concentrate on the scenario where the scalar field, as the result of contraction of the universe, is driven from the bottom of the potential well. The field is amplified, and finally the phase of the standard slow-roll inflation is realized. Such an evolution modifies the standard inflationary spectrum of perturbations by the additional oscillations and damping on the large scales. We extract the parameters of the model from the observations of the cosmic microwave background radiation. In particular, the value of inflaton mass is equal to  $m=(2.6 \\pm 0.6) \\cdot 10^{13}$ GeV. In our considerations we base on the seven years of observations made by the WMAP satellite. We propose the new observational consistency check for the phase of slow-roll inflation. We investigate the conditions which have to be fulfilled to make the observations of the Big Bounce effects possible. We translate them to the requirements on the parameters of the model and then put the observational constraints on the model. Based on assumption usually made in loop quantum cosmology, the Barbero-Immirzi parameter was shown to be constrained by $\\gamma<1100$ from the cosmological observations. We have compared the Big Bounce model with the standard Big Bang scenario and showed  that the present observational data is not informative enough to distinguish these models. }  \\begin{document} ",
        "introduction": "\\label{Intro}  The observations of the cosmic microwave background (CMB) radiation indicate that the power spectrum of primordial scalar perturbations is in the broad range \\emph{nearly} scale-invariant. Therefore, the spectrum can be written in the power-law form \\begin{equation} \\mathcal{P}_{\\text{s}}(k) = A_{\\text{s}} \\left( \\frac{k}{k_0}\\right)^{n_{\\text{s}}-1}, \\label{powerspectrum} \\end{equation} where the spectral index $n_{\\text{s}}$ is close to unity. Here, $A_{\\text{s}}$ is an amplitude of the scalar perturbations and $k_0$ is the so-called pivot number. The case $n_{\\text{s}}=1$ corresponds to the scale-invariant Harrison-Zeldovich spectrum. The observed spectrum is almost of this type, namely it is little red-shifted ($n_{\\text{s}} \\lesssim 1$). In particular, the seven years of observations made by the WMAP satellite \\cite{Komatsu:2010fb} indicate that $n_{\\text{s}} = 0.963\\pm 0.012$ (68\\% CL).  The spectrum in the form discussed above can be explained by the phase of cosmic inflation (see e.g. \\cite{Linde:2007fr}). This phase can be driven by the self-interacting scalar field, the so-called inflaton field. In the most conservative approach the inflation can be driven by a single massive scalar field. This case will be considered in this paper. In this model, a nearly scale-invariant spectrum spectrum is generated during the slow-roll phase. In a more general case the inflation can be driven with the different potentials. However, the potentials other than massive, lead to the non-Gaussian structure of the cosmic primordial perturbations what can be constrained by the CMB observations \\cite{Bartolo:2004if}. Since the CMB anisotropies do not indicate any non-Gaussian signatures, the massive potential is somehow privileged. However, with the present sensitivity on these kind of effects, some of the other potentials are still allowed.  Other models, as multi-field inflation are also possible to be realized. In this paper, we consider the simplest possible realization of the inflation which is given by the single massive scalar field.  The weak point of the slow-roll inflationary scenario is that it requires some special initial conditions. Namely, the field has to start its evolution not from the bottom of potential well but from the position which is far from its center. In the classical model there is no mechanism to drive this field up the potential well. However, it has been recently pointed out \\cite{Mielczarek:2009zw} that the phase of a quantum bounce can drive the inflaton field up the potential well and set the proper initial conditions for the slow-roll phase. The studies were performed within the loop quantum cosmology (LQC) \\cite{Bojowald:2008zzb,Ashtekar:2008zu}, however the mechanism is generic for all models with the bouncing phase. In the framework of LQC this issue has been studied recently in \\cite{Mielczarek:2009zw,Ashtekar:2009mm,Chiou:2010nd,Mielczarek:2010bh}.  In the framework of LQC the classical dynamics of the universe is significantly modified when the energy density approaches the Planck energy density. These effects of the quantum gravitational modification can be introduced as corrections to the classical equations of motion. In particular, the modified Friedmann equation takes the form \\begin{equation} \\left(\\frac{1}{a}\\frac{da}{dt}\\right)^2= \\frac{8\\pi}{3m^2_{\\text{Pl}}} \\rho\\left(1-\\frac{\\rho}{\\rho_{\\text{c}}} \\right), \\label{Friedmann} \\end{equation} where the critical energy density is defined as follows \\begin{equation} \\rho_{\\text{c}}=\\frac{\\sqrt{3}}{16\\pi^2\\gamma^3}\\rho_{\\text{Pl}}, \\label{rhoc} \\end{equation} where $\\rho_{\\text{Pl}}:=m^4_{\\text{Pl}}$ and $m_{\\text{Pl}}\\approx1.22\\cdot 10^{19}$ GeV is the Planck mass.  The above expression (\\ref{rhoc}) is crucial, because it relates the Barbero-Immirzi parameter $\\gamma$ with the parameter $\\rho_{\\text{c}}$, which can be constrained observationally. This is a good example of expressing the phenomenological parameter, as $\\rho_{\\text{c}}$, in terms of parameters of the underlying theory. In this case the fundamental theory is loop quantum gravity (LQG) with a free parameter $\\gamma$.  The value of Barbero-Immirzi  parameter is usually fixed from the considerations of the black hole entropy \\cite{Ashtekar:1997yu}. In particular, the value $\\gamma=0.239$ \\cite{Meissner:2004ju} was derived what leads to $\\rho_{\\text{c}}=0.82 \\ m^4_{\\text{Pl}}$. The expression (\\ref{rhoc}) is however not free from ambiguities. In particular, it bases on assumption that the area of the loop in LQC is equal to a gap of the area operator within LQG. This is not necessary true and therefore findings based on expression (\\ref{rhoc}) must be carried with a due care. In general, $\\rho_{\\text{c}}$ can be treated as a free phenomenological parameter (see e.g. \\cite{Dzierzak:2008dy,Malkiewicz:2009zd, Mielczarek:2010rq}). ",
        "conclusions": "\\label{Summary}  In this paper we have examined observations of the cosmic microwave background radiation as the potential probe of physics in the Planck epoch. We have based our considerations on the predictions of loop quantum cosmology.  Within this approach the initial singularity is replaced by a cosmic bounce.  During the bounce, the universe reaches the maximal energy density $\\rho_{\\text{c}} < \\infty$. In our studies, we have concentrated on the model with a massive scalar field. The advantage of this choice is that the phase of the standard slow-roll inflation is realized during the evolution. During this phase, the primordial perturbations are formed. Their properties can be investigated by observations of the CMB. Based on the recent results from the seven years of observations made by the WMAP satellite, we have determined some parameters of the model. In particular, we have found that inflaton mass $m=(1.4 \\pm 0.5) \\cdot 10^{-6} m_{\\text{Pl}}=(2.6 \\pm 0.6) \\cdot 10^{13}$ GeV.  Subsequently, we have investigated the modifications  of the primordial power spectrum due to presence of the bounce.  The two main effects that were discussed are: suppression and oscillations of the inflationary spectrum. The suppression can very strongly modify the spectrum. However, the present scale of suppression was shown to be behind the Hubble radius. Therefore, the effect of suppression is not directly observationally available.  Despite this, the present scale of suppression $\\lambda_*=\\frac{2\\pi}{k_*}  \\approx 40 \\ \\text{Gpc}$ was distinguished by the observations. This is because of the oscillations in the primordial power spectrum. The effect of oscillations is in general, more subtle and dominated by the cosmic variance. However, the present work indicates that the oscillations in the power spectrum can in fact explain the strong additional modulations in the spectrum of CMB on the low multipoles. In order to verify it, we have compared the Big Bounce model with the standard Big Bang scenario and showed  that the present observational data is not informative enough to distinguish these models. In other words, the Big Bounce predictions were shown not to be  in conflict with the observational data.  Moreover, despite the Big Bounce model has one more parameter $k_*$, the evidence obtained in this case is comparable with the Big Bang case.  We have shown that $\\phi_{\\text{max}}$ is not lower  than $2.97 m_{\\text{Pl}}$. Based on this, some models of the bounce, as the mentioned \\emph{symmetric} model, can be  significantly constrained or even excluded. Assuming validity of (\\ref{rhoc}) we gave an observational constraint on the Barbero-Immirzi parameter. We found that $\\gamma < 1100$. The corresponding constraint on $\\rho_{\\text{c}}$ was shown to be much weaker."
    },
    "1005/1005.4356_arXiv.txt": {
        "abstract": "The luminosity, light curve, post--maximum spectrum, and lack of a progenitor on deep pre-outburst images suggest that SN 2010U was a luminous, fast nova. Its outburst magnitude is consistent with that  for a fast nova using the Maximum Magnitude-Rate of Decline relationship for classical novae. ",
        "introduction": "Episodic high mass loss events dominate the final stages of massive star evolution. The evidence for these events is observed across the upper HR diagram ranging from giant eruptions like $\\eta$ Car, and the ``LBV nebulae'' associated with many of the Luminous Blue Variables to the cool hypergiants and red supergiants like IRC~+10420 and VY CMa with their complex circumstellar ejecta.  Unfortunately, the observational record is sparse because these stars are rare and their importance has only been fully recognized in recent years.   Modern supernova surveys are finding  a growing assortment of similar and related objects. These ``impostors'' are under-luminous, have spectra with narrow emission lines and much slower ejection velocities. Some of these objects may be undergoing `giant' eruptions possibly similar to $\\eta$ Car but with much shorter durations (e.g., SN 2009ip, \\citet{Smi2010}). Others appear to be generically related to the normal LBV/S Dor variables like SN 2002kg (Var 37 in NGC 2403; \\citet{KW05,vandyk06,Maund06}). A third subgroup including SN 2008S and the NGC 300 optical transient were heavily obscured prior to the outburst \\citep{JPetal08, JP08}; their progenitors  may have recently evolved from a high mass losing stage such as an AGB star or red supergiant \\citep{Smi09,Bond09,Berg09,Bott09,Thomp09,JPetal09}. We used to think that ``eruptions''  in which the star increased its total luminosity were primarily associated with the most massive stars, upwards of 40 --50 M$\\odot$, but the latter pair apparently originated from 10 -- 20 M$\\odot$ stars (references above and \\citet{Gogarten}). It is thus increasingly apparent that these eruptions or outbursts may not all be the same phenomenon; the objects  represent a range of stellar masses and may originate from stars in different evolutionary states.  Although discovery of these transient non-terminal eruptions are becoming more common, our information  about them is still very incomplete. The most recent addition, SN2010U in NGC 4214,  was first reported in eruption on 5 February, 2010 (S. Nakano, CBET 2161) at apparent magnitude $\\sim$ 16 on unfiltered CCD frames.  An early spectrum (Marion, Vinko, \\& Wheeler, CBET 2163) with strong, narrow emission lines  of hydrogen, Ca II and Na I with P Cyg absorptions, quickly showed that it was not a true supernova, but possibly an eruptive variable or transient.  In this {\\it Letter} we  present pre and post-eruption observations of SN 2010U and a analysis of the associated stellar population in NGC 4214. Our discussion of its light curve, spectrum,  and  limitations on  the mass and luminosity of its likely progenitor suggest that SN2010U is a very luminous nova, not the eruption of a massive star. ",
        "conclusions": "The rapid decline of the light curve, the post-maximum spectrum, and the relatively low mass and luminosity inferred from the lack of a precursor on deep pre-eruption images and the associated stellar population, all suggest that SN 2010U was a luminous nova.  SN 2010U shares characteristics with the fast novae. If we adopt the observations  by  the amateur observers and assume that it was at or near maximum, then the time to decline by two magnitudes (t$_{2}$) is 15 days and the time to decline three magnitudes (t$_{3}$) is 26 days. These decline rates are consistent with a  fast nova. Classical novae obey the Maximum-Magnitude-Rate of Decline relation (MMRD) of \\cite{Della95}. Adopting M$_{r}$, of $-10.5$  for maximum light and a $V-R$ color of 1.1 to 1.5 mag from V1500 Cyg, a classical nova a few days past maximum light \\citep{Jay76}, M$_{v}$ is $\\approx$ $-9$  to $-9.4$~mag, near the upper luminosities observed for CNe. A t$_{2}$ of 15 days implies a maximum M$_{v}$ of $-8.7$ mag \\citep{DD2000}, comparable to  SN 2010U's maximum luminosity. Of course, it is possible that maximum was missed and SN 2010U was  more luminous, but the P Cyg profiles reported in CBET 2163 indicate that it was not missed by more than a few days. Furthermore, a typical fast to moderately fast nova, with t$_{2}$ $\\sim$ 15 days, has emission lines with FWHM from  2000 km s$^{-1}$ up to as high as 6000 km s$^{-1}$. SN 2010U is on the low end but consistent with  this range. \\citet{Kas2010} have recently described a group of fast and luminous novae that are apparently inconsistent with the MMRD relation. However, SN 2010U is at the upper end of luminosities expected from the MMRD relation for normal novae and was super-Eddington at maximum.   While not a high mass eruptive variable like most of the other ``SN impostors'', SN 2010U was indeed an impostor and is an excellent example of the need for post--maximum photometry and spectra, and a study of the stellar environment to determine the nature and evolutionary state of these objects."
    },
    "1005/1005.0575_arXiv.txt": {
        "abstract": "We report spectroscopic observations with the Gemini South Telescope of 5 faint (V$\\sim 20$) RR Lyrae stars associated with the Pisces overdensity. At a heliocentric and galactocentric distance of $\\sim80$ kpc, this is the most distant substructure in the Galactic halo known to date. We combined our observations with literature data and confirmed that the substructure is composed of two different kinematic groups. The main group contains 8 stars and has $\\langle V_{gsr}\\rangle=50$ \\kms, while the second group contains four stars at a velocity of $\\langle V_{gsr}\\rangle=-52$ \\kms, where $V_{gsr}$ is the radial velocity in the galactocentric standard of rest. The metallicity distribution of RR Lyrae stars in the Pisces overdensity is centered on $[Fe/H]=-1.5$ dex and has a width of 0.3 dex. The new data allowed us to establish that both groups are spatially extended making it very unlikely that they are bound systems, and are more likely to be debris of a tidally disrupted galaxy or galaxies. Due to small sky coverage, it is still unclear whether these groups have the same or different progenitors. ",
        "introduction": "}  In the last decade, abundant observational evidence has been accumulated in favor of a hierarchical formation scenario of the halo of the Milky Way (see a review by \\citealt{helmi08}). The substructures in the halo have been observed both as overdensities in space and as moving groups. The tidal streams of the disrupting Sagittarius dwarf spheroidal (dSph) galaxy are the best example of such substructures, with the streams wrapping around most of the sky \\citep{ivezic00,yanny00,vivas01,ibata02,majewski03}. Other known substructures include the Virgo Stellar Stream \\citep{duffau06,prior09} and several other small structures associated with the Virgo overdensity \\citep{newberg07,juric08,vivas08}, the Monoceros structure \\citep{ibata03,yanny03,ivezic08}, the Triangulum-Andromeda overdensity \\citep{rocha04}, the Hercules-Aquila cloud \\citep{belokurov07a}, the Orphan stream \\citep{belokurov07b}, and other smaller overdensities \\citep{clewley06,starkenburg09} and streams \\citep{grillmair09,schlaufman09}. Using Sloan Digital Sky Survey (SDSS; \\citealt{york00}) multi-epoch data in the stripe 82 region ($308\\arcdeg<R.A. <60\\arcdeg$, $|Dec|<1.25\\arcdeg$), \\citet{sesar07} found an overdensity of RR Lyrae stars in the halo at a heliocentric and galactocentric distance of $\\sim80$ kpc and $R.A.\\sim354\\arcdeg$ (termed ``Structure J'', see Figures 10 and 11 from \\citealt{sesar10}). This structure was subsequently renamed the ``Pisces overdensity'' by \\citet{watkins09} after the constellation where it is located.  Convincing evidence that the Pisces overdensity is indeed a coherent structure, and not a random superposition of stars, was recently presented by \\citet{kollmeier09}. They measured radial velocities of 8 RR Lyrae stars located in the Pisces overdensity and found that 5 of them have a narrow velocity distribution centered on $\\sim -75$ {\\kms } in the heliocentric rest frame. The authors found hints of a second velocity peak in their data, containing two or maybe three stars, at a lower velocity ($<-150$ {\\kms } in the heliocentric rest frame). \\citet{kollmeier09} interpreted the main peak as a dwarf galaxy or a disrupting dwarf galaxy, and did not consider the second peak as a real group due to a small number of stars. In this paper, we expand on their work by reporting radial velocities and metallicities for 5 RR Lyrae stars in the Pisces overdensity, four of which were not observed by \\citet{kollmeier09}. By combining these two data sets, we are now in position to confirm the existence of a second peak in the velocity distribution, and are able to better explore the spatial extent and metallicity distributions of these two substructures present in the Pisces overdensity. ",
        "conclusions": "We confirm the existence of two kinematic groups in the Pisces overdensity, located at 83 kpc from the Sun, by combining new spectroscopic observations with those previously obtained by \\citet{kollmeier09}. The spatial extent of groups in the right ascension direction suggests that these are not bound systems but rather debris of disrupted dwarf galaxies, most likely tidal streams. If these are tidal streams, the question remains whether they have the same or different progenitors. The first case (same progenitor) is not an improbable one. The Sagittarius dSph galaxy has already shown us examples where debris stripped in different epochs overlap in the same region of the sky (see Figure 1 in \\citealt{law05}). In such regions large gradients in $V_{gsr}$ are evident along the same line of sight as some streams recede while others approach. The second case (different progenitors) is supported by state-of-the-art simulations of galaxy formation. These simulations predict numerous substructures of different morphologies (streams, shells, clouds) in halos of Milky Way-size galaxies \\citep{johnston08}. Such substructures may overlap and may be difficult to separate without wide and deep sky coverage.  Ongoing and upcoming wide-area surveys, such as Pan-STARRS \\citep{kaiser02} and LSST \\citep{ivezic08b}, will provide the depth and sky coverage needed to separate different halo substructures in coordinate space. With single-epoch and co-added observations, these surveys will be able to observe main sequence F dwarfs up to distances of 100 kpc and beyond, and will trace halo substructures with greater fidelity. In addition, detections of overdensities in main sequence and RR Lyrae stars will allow robust measurements of metallicities using a photometric method described in \\citet{sesar10}."
    },
    "1005/1005.5166_arXiv.txt": {
        "abstract": "T Pyx is a luminous recurrent nova that accretes at a much higher rate than is expected for its photometrically determined orbital period of about 1.8 hours. We here provide the first spectroscopic confirmation of the orbital period, $P =1.8295$ hours ($f = 13.118368 \\pm 1.1 \\times 10^{-5}$ c/d), based on time-resolved optical spectroscopy obtained at the VLT and the Magellan telescopes. We also derive an upper limit of the velocity semi-amplitude of the white dwarf,  $K1 = 17.9 \\pm 1.6 \\,$km$\\,$s$^{-1}$, and estimate a mass ratio of $q = 0.20 \\pm 0.03$. If the mass of the donor star is estimated using the period-density relation and theoretical main-sequence mass-radius relation for a slightly inflated donor star, we find $M_{2} = 0.14 \\pm 0.03 \\,$ M$_{\\odot}$. This implies a mass of the primary white dwarf of $M_{1} = 0.7 \\pm 0.2 \\,$ M$_{\\odot}$. If the white-dwarf mass is $>$ 1 M$_{\\odot}$, as classical nova models imply, the donor mass must be even higher. We therefore rule out the possibility that T Pyx has evolved beyond the period minimum for cataclysmic variables. We find that the system inclination is constrained to be $i \\approx 10$ degrees, confirming the expectation that T Pyx is a low-inclination system. We also discuss some of the evolutionary implications of the emerging physical picture of T Pyx. In particular, we show that epochs of enhanced mass transfer (like the present) may accelerate or even dominate the overall evolution of the system, even if they are relatively short-lived. We also point out that such phases may be relevant to the evolution of cataclysmic variables more generally. ",
        "introduction": "Binary systems containing a Roche-lobe filling main-sequence star that loses mass to a primary white dwarf (WD) are referred to as cataclysmic variables (CVs, see~\\citealp{b28} for an overview). T Pyx is a luminous CV in which the donor star is transferring mass at a very high rate onto a high-mass white dwarf, resulting in unusually frequent thermo-nuclear runaways (TNRs) on the surface of the primary. Between the years 1890 $-$ 1967, T Pyx has undergone five such nova eruptions, with an recurrence time of about 20 years between the eruptions, and was therefore classified as a member of the recurrent nova (RN) subclass. However, the last eruption was in 1966, which means that T Pyx has now passed its mean recurrence time by more than 20 years. The eruptive behavior of RNe in comparison with classical novae is thought to be due to a high mass-transfer rate in combination with a massive primary white dwarf ~\\citep{b31}.  ~\\cite{b13} carried out an extensive photometric study of T Pyx and found a stable, periodic signal at $P=1.83$ hours that was interpreted as the likely orbital period. This would place T Pyx below the CV period gap and suggests a donor mass around $M_{2} \\sim 0.1$ M$_{\\odot}$. This is surprising. According to standard evolutionary models, a CV below the period gap should be faint and have a low accretion rate driven primarily by gravitational radiation (GR). Yet T Pyx's quiescent luminosity and status as a RN both imply that it has a high accretion rate of $> 10^{-8}$ M$_{\\odot}$ yr$^{-1}$ (\\citealp{b13};~\\citealp{b20}). This is about two orders of magnitude higher than expected for ordinary CVs at this period.  Assuming that the determination of the photometric orbital period is correct, the existence of T Pyx is interesting for at least two reasons. First, unless we are seeing the system in a transient evolutionary state, its lifetime would have to be very short $\\tau \\sim M_{2}/\\dot{M}_{2} \\leappeq 10$~Myrs. This would imply the existence of an evolutionary channel leading to the fast destruction of at least some short-period CVs. Second, TNRs frequent enough to qualify as RNe are thought to be possible only on high-mass accreting WDs ($M_{1} \\gtappeq 1$ M$_{\\odot}$). Moreover, RNe are the only class of novae in which the WD is expected to gain more mass between eruptions than it loses during them. This would make T Pyx a strong candidate Type Ia supernova progenitor.  However, the recent study of the system by~\\cite{b18} (see also~\\citealp{b20}) suggests, first, that T Pyx {\\em is}, in fact, in a transient evolutionary state, and, second, that, integrated over many nova eruptions, its WD does lose more mass than it gains. More specifically,~\\cite{b18} suggest that T Pyx was an ordinary cataclysmic variable until it erupted as a nova in 1866. This eruption triggered a wind-driven supersoft X-ray phase (as first suggested by~\\citealp{b7}), resulting in an unusually high luminosity and accretion rate. However, unlike in the original scenario proposed by~\\cite{b7}, the supersoft phase is not self-sustaining, so that the accretion rate has been declining ever since the 1866 nova eruption from $\\dot{M} \\sim 10^{-7}$ M$_{\\odot}$\\,yr$^{-1}$ to $10^{-8} $M$_{\\odot}$\\,yr$^{-1}$. As a result, T Pyx has faded by almost 2 magnitudes since the nova eruption (\\citealp{b18}). Based on this, and the fact that T Pyx has already passed its mean recurrence time by more than 20 years,~\\cite{b18} argue that T Pyx might no longer even be a recurrent nova. If these ideas are correct, T Pyx is not a viable SN Ia progenitor, and its remaining lifetime can be substantially longer than a few million years. However, if all its ordinary nova eruptions are followed by relatively long-lived (> 100 yrs) intervals of wind-driven evolution at high $\\dot{M}$, its secular evolution may nevertheless be strongly affected, with significant implications for CV evolution more generally (see also Section~\\ref{SD}).  A key assumption in virtually all of these arguments is that the photometric period  measured by~\\cite{b13} is, in fact, the orbital period of the system. So far, there has only been one attempt to obtain a spectroscopic period for T Pyx, by~\\cite{b26}, who reported a spectroscopic modulation with $P=3.44$ hours. Such a long orbital period above the CV period gap would be much more consistent with the high accretion rate found in T Pyx. In this study, we present the first definitive spectroscopic determination of the orbital period of T Pyx, showing that it is, in fact, consistent with Patterson et al.'s photometric period. We also use our time-resolved spectroscopy to estimate the main system parameters, such as the velocity semi-amplitude of the white dwarf ($K1$), the mass-ratio (q), the masses ($M_{1}$ and $M_{2}$) and the orbital inclination ($i$). Finally, we discuss the implications of our results for the evolution of T Pyx and related systems. ",
        "conclusions": "\\label{SD}  The main result of our study is the spectroscopic determination of T Pyx's orbital period, $P_{orb} \\simeq 1.83$~hrs. This confirms that the system is a CV below the period gap and implies that its current accretion rate is at least 2 orders of magnitude higher than that of an ordinary CV at this period. We also find that our spectroscopic orbital period is consistent with the photometric ephemeris found for T Pyx (\\citealp{b13}, an updated version is given in Section~\\ref{PE}). This means not only that photometric timings can be used as a more precise and convenient tracer of the orbital motion, but also that the large period derivative required by the photometric ephemeris marks a genuine change in the orbital period of the system. In fact, the spectroscopic data are consistent with the photometric period only if the period derivative is accounted for. The period derivative obtained here from the latest combined photometric data ($\\dot{P} = 6.7 \\times 10^{-10}$) is slightly lower than that obtained by~\\cite{b13} from data up to 1997 ($\\dot{P} \\simeq 9 \\times 10^{-10}$). A decrease in the rate of period change would be in line with~\\cite{b18} scenario that T Pyx's days as a high-$\\dot{M}$ recurrent nova are numbered (at least until its next ordinary nova eruption). In any case, the average time-scale for period change found across all the photometric data are about $3 \\times 10^5$~yrs.  We have also used our spectroscopic data to obtain estimates of other key system parameters, most notably the radial velocity semi-amplitude of the WD, $K1 = 17.9 \\pm 1.6$ km\\,s$^{-1}$, and the mass ratio, $q = 0.20 \\pm 0.03$. The latter estimate rests on three key assumptions: first, that our determination of $K1$ is correct, second, that our estimate $v(R_{disk})\\sin i$ is correct, and, third, that the accretion disk around the WD in T Pyx is tidally limited. Taken at face value, this relatively high value of the mass ratio implies that the donor star in the system is not a brown dwarf. Thus T Pyx is not a period bouncer. If we assume that the radius of the secondary is $10\\% \\pm 10\\%$ inflated relative to an ordinary main sequence star of the same mass, we find that its most likely mass is $M_{2} = 0.14 \\pm 0.03$ M$_{\\odot}$.  Overall, the physical picture that emerges from our study is consistent with the scenario proposed by~\\cite{b18}. In particular, they suggest that, prior to the 1866 eruption, T Pyx was an ordinary CV. That eruption then triggered a  high-$\\dot{M}$ wind-driven phase, as suggested by~\\cite{b7} to account for T Pyx's exceptional luminosity. However, this phase is not quite self-sustaining, so that T Pyx is now fading and perhaps not even a RN anymore. In line with this picture, we find that the mass ratio and donor mass we derive are not abnormally low for a CV at its orbital period. This shows that the present phase of high-$\\dot{M}$ accretion cannot have gone on for too long already. The hint of a declining period derivative may point in the same direction, but this needs to be confirmed.  Does all this mean that the phase of extraordinarily high accretion rates T Pyx is currently experiencing will have no lasting impact on its evolution? Not necessarily. In a stationary Roche-lobe-filling system, the orbital period derivative and total mass-loss rate from the donor (wind loss + mass transfer) are related via \\begin{equation} \\frac{\\dot{P}_{orb}}{P_{orb}} = \\frac{3 \\zeta -1}{2} \\frac{\\dot{M}_2}{M_2} \\end{equation} where $\\zeta$ is the donor's mass-radius index. In T Pyx, which has an increasing orbital period, we will take $\\zeta \\simeq -1/3$, which is appropriate for adiabatic mass-loss from a fully convective star (e.g.~\\citealp{b7}). We thus expect that $\\dot{M_2}/M_2 \\simeq {\\dot{P}_{orb}}/{P_{orb}}$, which suggests a typical mass-loss rate from the donor of $\\dot{M_2} \\sim 5\\times10^{-7}$~M$_{\\odot}$/yr in its current state. (The accretion rate onto the WD can be lower than this, since, in the wind-driven scenario, much of this mass escapes in the form of an irradiation-driven outflow from the donor.)  If every ordinary nova eruption in T Pyx is followed by $\\sim$ 100 yrs of such high-$\\dot{M}_2$ evolution, the total mass loss from the donor in the luminous phase is $\\sim 5 \\times 10^{-5}$ M$_{\\odot}$ between any two such eruptions. This needs to be compared to the mass lost from the donor during the remaining part of the cycle. If this is driven by gravitational radiation, the mass loss rate from the donor will be $\\dot{M}_2 \\simeq 5 \\times 10^{-11}$~M$_{\\odot}$/yr~\\citep{b7}. The recurrence time of ordinary nova eruptions for such a system is on the order of $10^5$ yrs~\\citep{b31}, so the total mass lost from the donor during its normal evolution (outside the wind-driven phase) is $5 \\times 10^{-6}$ M$_{\\odot}$. {\\em This shows that the long-term secular evolution may be dominated by its high-$\\dot{M}$ wind-driven phases, even if the duty cycle of these phases is very low (e.g. 0.1\\% for the numbers adopted above).}  It is finally tempting to speculate on the relevance of ``T Pyx-like'' evolution for ordinary CVs. At first sight, the numbers above suggest that the evolution of a CV caught in such a state may be accelerated by about an order of magnitude. This is interesting, since it could bear on the long-standing problem that there are fewer short-period CVs and period bouncers in current samples than theoretically expected (e.g.~\\citealp{b13};~\\citealp{b14};~\\citealp{b15}; Pretorius \\& Knigge 2008ab). Moreover, this channel need not be limited to systems containing high-mass WDs that are capable of becoming RNe. After all, it is not the recurrent nova outbursts that are of interest from an evolutionary point of view, but simply the existence of a prolonged high-$\\dot{M}$ phase in the aftermath of nova eruptions. In CVs containing lower-mass WDs, such a phase may still occur, although its evolutionary significance could still depend on the WD mass. For example, the interoutburst time-scale is longer for low-mass WDs, and the duration of the high-$\\dot{M}$ phase could also scale with WD mass. One obvious objection to this idea is that, observationally, most nova eruptions are not followed by centuries- (or even decades-) long high-$\\dot{M}$ phases. However, this need not be a serious issue. Most observed novae are long-period systems, so if the triggering of a wind-driven phase requires a fully convective donor, most novae would not be expected to enter such a phase. It may be relevant in this context that at least one other short-period nova -- GQ Mus -- exhibited an exceptionally long post-outburst supersoft phase of $\\sim 10$~years~\\citep{b12}.~\\cite{b3} have also suggested that the duration of the supersoft phase in novae may scale inversely with orbital period.  However, there is another important consequence to the idea that the evolution of many short-period CVs is accelerated by ``T Pyx-like'' high-luminosity phases. The orbital period-derivative is positive in the high-$\\dot{M}$ phase, but negative during the remaining times of GR-driven evolution. But since $\\dot{P}_{orb}/P_{orb} \\simeq \\dot{M}_2/M_2$ in {\\em both} phases, the sign of the secular (long-term-average) period derivative will generally correspond to the phase that dominates the secular evolution. So whenever wind-driving dramatically accelerates the binary evolution, the direction of period evolution will also be reversed. Is this a problem? Perhaps. The recent detection of the long-sought {\\em period spike} in the distribution of CV orbital periods~\\citep{b4} suggests that there is, in fact, a reasonably well-defined minimum period for CVs, as has long been predicted by the standard model of CV evolution (e.g.~\\citealp{b9}). On the other hand, the location of the observed spike ($\\simeq 83$~mins) is significantly different from the expected one ($\\simeq 65-70$~min). Is it possible that the observed period minimum corresponds to the onset of wind-driving in most CVs? T Pyx, with $P_{orb} \\simeq 110$~min would then have to be an outlier, however, perhaps because of an unusually high WD mass.  The idea that T Pyx-like phases may significantly affect the evolution of many CVs is, of course, highly speculative, and we do not mean to endorse it too strongly. However, it highlights the importance of understanding T Pyx: until we know what triggered the current high-luminosity state, it will remain difficult to assess the broader evolutionary significance of this phase. Note that the apparent uniqueness of T Pyx is not a strong argument against such significance. For example, if the duty cycle of high-luminosity phases is $\\sim 0.1$\\%, as suggested by the numbers above, we should not expect to catch many CVs in this state. Thus T Pyx could be the tip of the proverbial iceberg."
    },
    "1005/1005.3377_arXiv.txt": {
        "abstract": "{ The issue of the influence of coronal holes (CHs) on coronal mass ejections (CMEs) in causing solar energetic particle (SEP) events is revisited. It is a continuation and extension of our previous work \\citep{Shen_etal_2006}, in which no evident effect of CHs on CMEs in generating SEPs were found by statistically investigating 56 CME events. This result is consistent with the conclusion obtained by Kahler in 2004. In this paper, we extrapolate the coronal magnetic field, define CHs as the regions consisting of only open magnetic field lines and perform a similar analysis on this issue for totally 76 events by extending the study interval to the end of 2008. Three key parameters, CH proximity, CH area and CH relative position, are involved in the analysis. The new result confirms the previous conclusion that CHs did not show any evident effect on CMEs in causing SEP events. ",
        "introduction": "Gradual solar energetic particle (SEP) events are thought to be a consequence of CME-driven shocks generating plenty of SEPs which would be observed near the Earth. In our previous work in 2006, we statistically studied the effect of coronal holes (CHs) on the CMEs causing SEP events by investigating the CME source locations and their relation with the CHs identified in EUV 284 \\AA$ $\\citep[][ hereafter Paper I]{Shen_etal_2006}. It was implied that neither CH proximity nor CH relative location exhibits any evident effect on the intensities of SEP events. This result is consistent with the conclusion obtained by \\citet{Kahler_2004}, who comparatively studied the SEP events produced in the fast and slow solar wind streams and found no significant bias against SEP production in fast-wind regions which are believed to originate from CHs.  These findings seem not quite fit people's `common sense', because CHs are believed to be regions with low-density and low temperature in the corona \\citep[e.g.][]{Harvey_Recely_2002}, from which the solar wind is fast and the magnetic field is open, and therefor apparently three disadvantages for a CME to produce SEP may exist when it is near a coronal hole region. These advantage are: (1) the background solar wind speed $V_{sw}$ near CHs is larger than that in other regions; (2) the plasma density near CHs is much lower than that in other regions, so that the Alfv\\'{e}n speed $V_a$ is larger \\citep{Shen_etal_2007,Gopalswamy_etal_2008}; and (3) the magnetic field lines in CHs are open. The first two disadvantages suggest that a strong shock might be hardly produced near CHs. The third one implies that particles might be able to escape from the shock acceleration process earlier and easier. Thus, it can be expected that CHs would influence the CME in producing SEP events. The work by \\citet{Kunches_Zwickl_1999} was consistent with the picture depicted above. In their paper, they found that the CH may delay the onset times of SEPs when a CH is present between Sun-observer line and the solar source of the SEP event. They also speculate that the peak intensity could be influenced by CH. However, they did not statistical study such influence. It is hard to say that their conclusion is statistically significant.  In principle, CHs are open field regions, though they were first identified in observations \\citep[e.g.][]{Zirker_1977}. \\citet{Kunches_Zwickl_1999} identified CHs based on He 10830 \\AA$ $. In our 2006 work (Paper I), CHs were auto-determined based on  EUV 284\\AA$ $ images taken by SOHO/EIT. Thus, it is doubtable whether or not the CHs identified in EUV wavelengths really represent open field regions. Another doubt in our 2006 work is that only frontside CHs are taken into account. In order to remove the doubt and get a more reliable result, we look into this topic again by extrapolating coronal magnetic field instead of analyzing EUV images. The term `CHs' in this paper therefore actually refers to open field regions. The magnetic field extrapolation and determination of CHs are introduced in Section 2. Section 3 presents the statistical analysis. A brief summary and conclusions are given in Section 4. ",
        "conclusions": "In order to study the influence of CHs on CMEs in producing SEP events,  a total of 76 west-side fast halo CMEs during 1997 - 2008 are investigated, as well as their associated CHs. Different from the CHs obtained by brightness method based on EIT 284\\AA$ $ data in paper I, the CHs we investigated in this paper are obtained with the aid of the extrapolation of coronal magnetic field by CSSS model, in which the MDI daily-updated synoptic magnetic field charts are adopted as the bottom boundary condition. By using this method, all the CHs, defined as the regions consisting of open magnetic field lines only, over the entire solar surface are inferred.  After analyzing three parameters, CH proximity, area of corresponding CHs and relative position between CHs and CMEs, it is found that all of the statistical results do NOT have significance exceeding the 1$\\sigma$ level. These parameters do NOT show any evident influence on SEP occurrence probability, and the speed of SEPYCMEs also do NOT show any difference between different groups binarized by these parameters.  These results confirmed the conclusion we got in Paper I and \\citet{Kahler_2004} that no evident influence of CHs on CME in producing SEP events.   An expanding CME may drive a quasi-parallel shock at its flank as discussed by \\citet{Kahler_2004}. The condition of CME in driven shock in this situation is $V_{cme}$ larger than local alf\\'{v}en speed $V_a$ or sound speed $C_s$ only. Thus, the fast flow speed near CHs may show no influence on producing strong shock. Beside, not only the plasma density but also the magnetic field strength in fast solar wind region is smaller than them in slow solar wind region\\citep{Ebert_etal_2009}, so the alf\\'{v}en speed in fast solar wind region may not obvious faster than it in slow solar wind region. Based on the these analysis, it could be expected that shock can also be produced in fast solar wind region near CH and no evident fast of the CME needed. In addition, the shock interact with background solar wind may generate a turbulence. Such turbulence could be treated as the main mechanism that makes particles back to shock acceleration process to produce SEP events\\citep{Reames_1999}. The close magnetic topology could only provide an addition method to make the particle back to shock acceleration\\citep{Shen_etal_2008}. So, the influence of open magnetic field topology may weak in shock producing SEP events.  \\normalem"
    },
    "1005/1005.4557_arXiv.txt": {
        "abstract": "{While sub-micron- and micron-sized dust grains are generally well mixed with the gas  phase in protoplanetary disks, larger grains will be partially decoupled and as a consequence have a different distribution from that of the gas.  This has ramifications for predictions of the observability of protoplanetary disks, for which gas-only studies will provide an inaccurate picture. Specifically, criteria for gap opening in the presence of a planet have generally been studied for the gas phase, whereas the situation can be quite different in the dust layer once grains reach mm sizes, which is what will be observed by ALMA.} {We aim to investigate the formation and structure of a planetary gap in the dust layer of a protoplanetary disk with an embedded planet.} {We perform 3D, gas$+$dust SPH simulations of a protoplanetary disk with a planet on a fixed circular orbit at 40 AU to study the evolution of both the gas and dust distributions and densities in the disk. We run a series of simulations in which the planet mass and the dust grain size varies.} {We show that the gap in the dust layer is more striking than in the gas phase and that it is deeper and wider for more massive planets as well as for larger grains. For a massive enough planet, we note that cm-sized grains remain inside the gap in corotation and that their population in the outer disk shows an asymmetric structure, a signature of disk-planet interactions even for a circular planetary orbit, which should be observable with ALMA.} {} ",
        "introduction": "\\label{sec:introduction}  While we understand the general scenario of planet formation via accretion \\citep{Mizuno1980, Wetherill1980, Lissauer1987, Wetherill1990, Pollack_etal1996}, the devil is in the detail and the processes by which tiny sub-micron grains grow into planetesimals, the building blocks of planets, is not well understood. With over 400 extrasolar planets detected to date as listed in the extrasolar planet encyclopedia\\footnote{http://exoplanet.eu/}, we can start to do some meaningful statistics on the types of planets that exist in our galaxy \\citep{AMBW05,M07} to help constrain theories of planet formation.  A wealth of analytical as well as numerical studies of the formation of gaps by a planet in a gas disk have shown how the shape of the gap depends on properties of the planet (mass) as well as the gaseous disk (pressure scale height, viscosity). See \\cite{CMM06} for the case of gap shapes and \\cite{PNKMA07} for a more general review on planet migration and gap formation.  The dust phase, however, has been shown to behave differently to the gas. Dust experiences a headwind from the pressure-supported sub-Keplerian gas and the induced drag force slows the dust and makes it settle to the midplane and migrate inwards. The magnitude of these effects depends strongly on the grain size and the disk density \\citep{W77,SV96,SV97,GBFL04,GL04,BF05}.  In recent years, gap formation by a planet embedded in disks of gas and dust has been studied by several authors, for different planet masses and different sizes of the solid particles. \\citet{PM04,PM06} determined the spatial distribution of 1~mm grains with 2D simulations in order to derive the smallest planet mass that would result in an observable gap with ALMA (Atacama Large Millimeter Array). In their work the dust was strongly coupled to the gas and responded indirectly to the planet gravity through the radial pressure gradients caused by a Neptune-mass planet in the gas disk that led to the formation of a gap. \\citet{Muto09} investigated the effect of a low-mass planet on the dust distribution by injecting one dust grain at a time and derived criteria for gap opening. \\citet{Ciecielag07} as well as \\citet{Marzari00} focused on already formed planetesimals while the gas phase is still present. \\citet{Ciecielag07} considered planetesimals down to 1~m in size. They studied the effect of spiral structures in the circumprimary gas disk triggered by a secondary companion in a tight binary system in order to derive relative velocities between planetesimals and determine whether planet formation is possible in such systems.  The presence of planets, and the gap they create when they are massive enough, can also help constrain the global properties of the gas (temperature, density, viscosity) as well as those of the dust (grain size distribution, degree of settling) in the disks. This can be achieved by measuring the width and brightness of the gap, ideally for each phase \\citep{W05,CMM06,Fouchet07}.  It has been shown that ALMA will be able to observe a planetary gap opened by a 1~$M_J$ planet at 5.2~AU from a 1~$M_\\odot$ star at a distance of 140~pc under the naive assumption that gas and dust are well mixed \\citep{W02}. It will certainly be possible to observe other features related to the presence of the planet, such as warm dust in its vicinity or even spiral waves for distances not exceeding 100~pc \\citep{W05}. Here we focus on the formation and features of the gap itself in the case of a massive protoplanet.  Our previous simulations of dust evolution in a typical Classical T-Tauri Star (CTTS) disk \\citep[hereafter BF05]{BF05}\\defcitealias{BF05}{BF05} showed that the thickness of the dust layer depends on grain size because different sized grains fall to the midplane at different rates. We distinguished three dynamical regimes for the dust: (1) almost uncoupled for large grains where the dust component follows slightly perturbed Keplerian orbits and keeps its 3D distribution (if initially 3D); (2) weakly coupled for intermediate-sized grains for which settling is very efficient; (3) strongly coupled for small grains where grains are forced to follow the gas motion.  We have also investigated the formation of a gap by a planet immersed in a Minimum Mass Solar Nebula (MMSN) and showed that dust settling makes the gap much more striking in the dust layer than in the gas phase because of its reduced vertical extension for weakly coupled grains \\citep[hereafter F07]{Maddison07,Fouchet07}\\defcitealias{Fouchet07}{F07}. Indeed, the criterion for gap formation depends on the disk scale height \\citep{CMM06}. Because of the size of particles in the weakly coupled regime for the particular case of the MMSN (1~m, see Sect.~\\ref{sec:theory}), that study had no direct application to observations. Images of dusty disks at infrared wavelengths do not probe such large particles, but instead trace the smaller, strongly coupled dust grains whose distribution is similar to that of the gas.  In this paper, we examine the gap formation in the dust layer of CTTS disks, which are spatially more extended and less dense than the theoretical MMSN case. Our ultimate goal is to use the results of our hydrodynamical simulations to produce synthetic images for ALMA. We consider in particular the effects of a massive planet in the outer cooler regions of the disk. Indeed, several planets at large distances from their star have been detected, such as those recently announced orbiting Fomalhaut \\citep{Kalas08}, or HR~8799 \\citep{Marois08}. The extrasolar planet encyclopedia lists 10 planets with a semi-major axis larger than 20~AU, and 7 of them have minimum masses larger than 5~$M_\\mathrm{J}$. We study the dust distribution in the weakly coupled regime, which corresponds for these CTTS disks to a size range (100 $\\mu$m to 1 cm) that can directly be probed by current and future (sub)millimetre instruments.  The present paper is the first part of this work, presenting the hydrodynamical simulations. In Sect.~\\ref{sec:theory}, we discuss the gas-dust interaction in the presence of a planet. In Sect.~\\ref{sec:simulations}, we describe the numerical method and simulation suite. Results are described in Sect.~\\ref{sec:results}, while the analysis and explanations are presented in Sect.~\\ref{sec:discussion}. We conclude in Sect.~\\ref{sec:conclusion}. In a forthcoming companion paper, we will use the simulations presented here to produce synthetic images for ALMA and present the most favorable observing configurations to detect the gap and associated structures. ",
        "conclusions": "\\label{sec:conclusion}  We have run 3D, two-phase (gas and dust) SPH simulations of a typical CTTS disk of mass 0.02~$M_{\\sun}$ with an embedded giant planet on a circular orbit at 40~AU. We vary the grain size (100~$\\mu$m, 1~mm, 1~cm) and the planet mass (0.1, 0.5, 1, 5~$M_\\mathrm{J}$) and study the formation of the planetary gap. We confirm that gap opening is stronger in the settled dust layer than in the flared gas disk. Gaps are deeper and wider for (1) larger, more efficiently settled grains and (2) more massive planets. Larger planet masses are required to open a gap in the gas phase than in the dust: while a 0.5~$M_\\mathrm{J}$ planet only slightly affects the gas phase, it carves a deep gap in the dust. For the most massive 5~$M_\\mathrm{J}$ planet, 1~cm grains remain trapped in corotation with the planet while their distribution in the outer disk shows an asymmetric structure, even though the planet's orbit is circular. We find that this is not caused by the periastron alignment of a coherent set of orbits but rather by the pile-up of dust in the pressure maximum of the gas phase caused by the spiral density wave triggered by the planet. This global asymmetry does not appear for less massive planets because the spiral perturbation is substantially weaker.  The variety of structures that we obtain in the dust phase for various grain sizes and planet masses has implications on the appearance of protoplanetary disks at (sub)millimetre wavelengths and show how important it is to go beyond the gas-only disk description proposed by, e.g., \\citet{W05}. The observability of these disks with ALMA is the subject of a forthcoming companion paper."
    },
    "1005/1005.5551_arXiv.txt": {
        "abstract": "{ Comprehensive VLBI and multi-waveband monitoring indicate that a single superluminal knot can cause a number of $\\gamma$-ray flares at different locations. However, the often very rapid variability timescale is a challenge to theoretical models when a given flare (perhaps the majority of those observed) is inferred from observations to lie near the 43 GHz core, parsecs from the central engine. We present some relevant observational results, using the BL Lac object AO~0235+164 as an example. We propose a turbulent cell model leading to a frequency-dependent filling factor of the emission region. This feature of the model can provide a solution to the timescale dilemma and other characteristics of blazar emission.   } ",
        "introduction": "There exists a crisis in the interpretation the results of multi-waveband monitoring of blazars. Comparisons of dates of peak $\\gamma$-ray flux measured by EGRET with epochs of ejections of superluminal radio knots (Jorstad et al.\\ \\cite{jor01a}), onsets of millimeter-wave flares (L\\\"ahteenm\\\"aki \\& Valtaoja \\cite{laht03}), and changes in centimeter-wave polarization (Jorstad et al.\\ \\cite{jor01b}) lead to the conclusion that most $\\gamma$-rays outbursts are coincident with radio events. Yet we know that the jets of blazars are opaque to radio emission within $\\sim 1$ pc or more from the central engine. This is difficult to reconcile with the shortest time scales of GeV variability observed --- a few hours in PKS~1622$-$297 (Mattox et al.\\ \\cite{mat97}) and 3C~454.3 (Foschini et al.\\ \\cite{fos10}).  However time scales of variability reflect size scales, not distance from the central engine. Furthermore, jets are very narrow, with opening half-angle $\\sim 10^\\circ/\\Gamma$, where $\\Gamma$ is the bulk Lorentz factor of the flow in the jet (Jorstad et al. \\cite{jor05}). Because of this, a jet can have a cross-sectional radius much less than its distance from the central engine. Nevertheless, this does not solve the entire problem. The core on 43 GHz Very Long Baseline Array (VLBA) images of PKS~1510$-$089 has been estimated to lie $\\sim 20$ pc from the central engine (Marscher et al.\\ \\cite{mar10a}). Offsets of parsecs are implied for other luminous blazars as well (e.g., Marscher et al.\\ \\cite{mar08}; Chatterjee et al.\\ \\cite{chat08}). The cross-sectional radius of a luminous blazar is $\\sim 0.1$ pc near the 43 GHz core. With a Doppler factor $\\delta \\sim 20$ and redshift $\\sim 0.5$, the shortest time scale of variability should be $\\sim 1$ week, not a fraction of a day.  Because of this discrepancy, a number of authors have insisted that the $\\gamma$-ray emitting region lies within $\\sim 10^{16}$ cm of the central engine (e.g., Tavecchio et al.\\ \\cite{tav10}). This has the advantage for high-energy emission models that optical-uv photons from the broad emission-line region are available for scattering to $\\gamma$-ray energies by highly relativistic electrons in the jet. Poutanen \\& Stern (\\cite{ps10}) find that pair production off the uv photon field can then cause the observed sharp break in the spectrum of some blazars at GeV energies. However, such an interpretation requires that the timing of the radio and $\\gamma$-ray events is a chance coincidence in every case. This seems unlikely, especially with singular events such as the ultra-high-amplitude optical flare, very sharp $\\gamma$-ray flare, and passage of a superluminal knot through the 43 GHz core in PKS~1510$-$089 at essentially the same time (Marscher et al.\\ \\cite{mar10a,mar10b}). Furthermore, the number of $\\gamma$-ray outbursts observed with {\\it Fermi} that are either coincident with or follow the passage of a new superluminal knot through the centroid of the core is becoming high enough to conclude that a large fraction of the $\\gamma$-ray emission originates parsecs away from the central engine (Jorstad et al. \\cite{jor10} and these proceedings).  We propose to solve this dilemma by developing a model in which much of the optical and high-energy radiation in a blazar is emitted near the 43 GHz core in VLBA images, parsecs from the central engine. The model allows for short time scales of optical and $\\gamma$-ray variability by restricting the highest-energy electrons radiating at these frequencies to small sub-regions of the jet. That is, in our model the filling factor at high frequencies is relatively low, while that of the electrons radiating at $\\sim 10^{10-13}$ Hz is near unity. Such a model is consistent with other prominent features of optical vs.\\ lower frequency emission. ",
        "conclusions": ""
    },
    "1005/1005.2237_arXiv.txt": {
        "abstract": "We investigate effects of hardness of primordial binaries on whole evolution of star clusters by means of $N$-body simulations. Using newly developed code, GORILLA, we simulated eleven $N=16384$ clusters with primordial binaries whose binding energies are equal in each cluster in range of $1-300kT_0$, where $1.5kT_0$ is average stellar kinetic energy at the initial time. We found that, in both soft ($\\le 3kT_0$) and hard ($\\ge 300kT_0$) limits, clusters experience deep core collapse. In the intermediate hardness ($10-100kT_0$), the core collapses halt halfway due to an energy releases of the primordial binaries. The core radii at the halt can be explained by their energy budget. ",
        "introduction": "Recently, observational informations concerning on the binary systems in the globular cluster have been accumulated from photometric observations of the eclipses (e.g. \\cite{Mateo96}), spectroscopic observations (e.g. \\cite{Albrow+01}), low-mass X-ray binaries (see \\cite{Liu+07}), and also the color-magnitude diagrams (e.g. \\cite{Rubenstein+97}). \\citet{Davis+08} also constrain the binary fraction in NGC 6397 by the method of the color-magnitude diagrams, and review the binary fractions in galactic globular clusters.  The binaries in the globular cluster play important roles for its dynamical evolution, since the binaries work as energy source through interactions with other stars or binaries (\\cite{Heggie75}). Even if the globular cluster has no binary at initial, the binaries are formed in due course through three-body encounter, and energy generated from these binaries halts core collapse of the globular cluster (\\cite{Henon75}).  Furthermore, if the globular cluster contains a nonnegligible fraction of primordial binaries, the dynamical evolution could be affected. \\citet{Goodman+89} (hereafter GH89) first showed theoretically difference of evolutions with and without primordial binary. GH89 estimated that the core of the cluster with primordial binaries at the halt of core contraction is larger than that of the cluster without primordial binaries by order of magnitudes. GH89's estimate is based on the model that the cluster core stop contracting when energy generated by the primordial binaries is balanced with energy outflowing from the inner region of the cluster to the outer region through two-body relaxation. McMillan et al. (1990; 1991) first performed $N$-body simulations of clusters with primordial binaries. They clearly showed that the cluster cores stop contracting at larger cores than those without primordial binaries. By means of Fokker-Planck model, \\citet{Gao+91} investigated post-collapse evolution of cluster with the primordial binary. They showed that the cores of the clusters continues to contract slowly after rapid core contraction and that the gravothermal oscillations occur after the several ten half-mass relaxation time of the slow core contraction.  In order to set a population of the primordial binaries in the cluster, several parameters concerning the binary and its distributions, such as mass fraction, and distribution of binding energies and eccentricities, need to be specified. Although these parameters should be derived from theories of star and cluster formations or be limited observationally, sufficient informations have not yet been provided at present. Only a few studies investigated the effect of these parameters on the evolution. \\citet{Heggie+92} performed $N$-body simulations of clusters with two different mass fractions, $6$ and $12$ \\%, of primordial binaries, and \\citet{Heggie+06} investigated the evolutions of clusters with $0-100$ \\% mass fraction of primordial binaries by means of $N$-body simulations. \\citet{Vesperini+94} (hereafter VC94) extended the model of GH89, taking into account the core mass fraction of primordial binaries and the distribution of the binding energies of the binaries.  In this paper, we focus on the dependence on binding energies of the primordial binaries. By means of $N$-body simulations, we systematically investigate the dynamical evolution of clusters with different initial hardness of binding energies. We set the distribution of binding energies of primordial binaries by delta function, $\\delta(x-E_{{\\rm bin},0})$, where $E_{{\\rm bin},0}$ is the initial binding energies of the primordial binaries and $E_{{\\rm bin},0}=1,3,10,30,100,300kT_0$. Here, $1.5kT_0$ is the average kinetic energy of stars in the cluster at the initial time. In previous simulations, the distribution of the binding energy of the primordial binaries are usually fixed at uniform distributions in $\\log E_{{\\rm bin},0}$ (see \\cite{McMillan+90}; \\cite{Gao+91}; \\cite{Heggie+92}; \\cite{Heggie+06}).  We found that the evolutions of the cores are different according to the distribution of the binding energies. If primordial binaries consist of softer binaries, the binaries are disrupted before core collapse through binary-single and binary-binary encounters, and do not affect the core evolution. If primordial binaries consist of harder binaries, the binaries escape from the cluster. This is because energy generated by the harder binaries per one encounter is larger than the cluster potential. They also do not affect the core evolution and the cluster exhibits deep core collapse. In intermediate range, the binaries efficiently heat the clusters, and the core collapses of the clusters halt halfway. These behaviors are consistent with the theoretical estimate by VC94.  These behaviors were not shown by previous works (\\cite{McMillan+90}; \\cite{Gao+91}; \\cite{Heggie+92}; \\cite{Heggie+06}; \\cite{Fregeau+07}; \\cite{Trenti+07}) who followed the dynamical evolution of clusters with primordial binaries. This is mainly because the previous works fixed the distributions of the binding energies of the primordial binaries at uniform distributions in $\\log E_{{\\rm bin},0}$, and their fixed distributions include binaries with hardness in soft, intermediate, and hard ranges.  The structure of this paper is as follows. We describe simulation methods in section \\ref{sec:sim}. In section \\ref{sec:results}, we present results of simulations and investigate the effect of the hardness of primordial binaries on the dynamical evolution of clusters. In section \\ref{sec:summary}, we summarize this paper. ",
        "conclusions": "\\label{sec:summary}  We study systematically the dependence of cluster evolution on the binding energy of primordial binaries. By means of GORILLA, we simulate the core evolution of the clusters, each of which contains primordial binaries with equal binding energy.  When the initial mass fraction of the primordial binaries is fixed to $0.1$, we find that the dynamical evolutions of the clusters are divided into three ranges according to hardness of the primordial binaries as follows. \\begin{enumerate} \\item In soft range ($<3kT_0$), the clusters experience core collapse in similar way to those without primordial binaries. The ratios of core radii to half-mass radii at the halt of the core collapse are about $0.006$. The primordial binaries do not heat the clusters. This is because the primordial binaries are destroyed through encounters with single stars, and do not generate energy. \\item In intermediate hard range ($10kT_0-100kT_0$), the core collapses in the clusters halt halfway. The ratios of core radii to half-mass radii at the halt of the core collapse are $0.05-0.1$. The primordial binaries release energy, and the energy heats the clusters. \\item In super hard range ($>300kT_0$), the clusters experience core collapse, and the ratios of core radii to half-mass radii at the halt of the core collapse is about $0.02$. The primordial binaries do not so much heat the clusters. Although the primordial binaries release energy through encounters, the energy is so large that binaries and single stars involved with the encounters are ejected from the clusters. \\end{enumerate}  The dependences of the boundaries between the soft and intermediate hard ranges and between the intermediate and super hard ranges on the initial mass fraction of the primordial binaries are as follows. \\begin{enumerate} \\item The boundary between the soft and intermediate hard ranges depends on the initial mass fraction of the primordial binaries. When the mass fraction of the primordial binaries is $0.3$, the core contraction in the cluster with $3kT_0$ primordial binaries halts at large ratio of core radius to half-mass radius $\\sim 0.07$, and the intermediate hard range includes $3kT_0$. \\item The boundary between the intermediate and super hard ranges is not changed, when the initial mass fraction of the primordial binaries ranges from $0.03$ to $0.3$. \\end{enumerate}  We compared the pairs of the ratios of core radii to half-mass radii and the core mass fraction of the binaries at the halt of the core contraction in our simulations with those of theoretical estimates. We found a good agreement between $N$-body simulations and the theoretical values."
    },
    "1005/1005.4231_arXiv.txt": {
        "abstract": "We describe in detail the general methodology and numerical implementation of consistent N-body simulations for coupled scalar field cosmological models, including the background cosmology and the generation of initial conditions (with the different couplings to different matter species taken into account). We perform fully consistent simulations for a class of coupled scalar field models with an inverse power-law potential and negative coupling constant, for which the chameleon mechanism does not operate. We find that in such cosmological models the scalar-field potential plays a negligible role except in the background expansion, and the fifth force that is produced is proportional to gravity in magnitude, justifying the use of a rescaled gravitational constant G in some earlier N-body simulations of similar models. We study the effects of the scalar coupling on the nonlinear matter power spectra and compare with linear perturbation calculations to investigate where the nonlinear model deviates from the linear approximation. For the first time, the algorithm to identify gravitationally virialized matter halos is adapted to the scalar field cosmology, and then used to measure the mass function and study the properties of virialized halos. We find that the net effect of the scalar coupling helps produce more heavy halos in our simulation boxes and suppresses the inner (but not the outer) density profile of halos compared with those predicted by lambda-CDM, while this suppression weakens as the coupling between the scalar field and dark matter particles increases in strength. ",
        "introduction": "\\label{sect:intro}  The nature of the dark energy \\cite{Copeland:2006} driving an apparent acceleration of the universe has been a cosmological puzzle for more than a decade. Amongst the models proposed to explain it, those incorporating scalar fields are by far the most popular, not only because of their mathematical simplicity and phenomenological richness, but also because the scalar field is a natural ingredient of many high-energy physics theories. A scalar field contributes a single dynamical degree of freedom which can interact indirectly with other matter species through gravity or couple directly to matter, producing a fifth force on the matter which creates violations of the Weak equivalence principle (WEP). This second possibility was introduced with the hope that such a coupling could potentially alleviate the coincidence problem of dark energy \\cite{Amendola:2000} and has since then attracted much attention (see, for example,\\emph{~}\\cite% {Bean:2001, Amendola:2004, Koivisto:2005, Boehmer:2008, Bean:2008, Bean:2008b, Boehmer:2009} and references therein for some recent work).  If there is a direct coupling between the scalar field and baryons, then the baryonic particles will experience a fifth force, which is severely constrained by observations, unless there is some special mechanism suppressing the fifth-force effects. This is the case in chameleon models, where the scalar field (the chameleon) gains mass in high-density regions (where observations and experiments are performed) and the fifth force effects are confined to undetectably small distances \\cite{Khoury:2004a, Khoury:2004b, Mota:2006, Mota:2007}. A common approach which avoids such complications is to assume that the scalar field couples only to the dark matter, an idea seen frequently in models with a coupled dark sector (\\emph{% e.g.}~\\cite{Maartens:2009, Valiviita:2010, Peacock:2010}). In this work our scalar field will not be chameleon-like as this case has been investigated elsewhere \\cite{Li:2009sy, Zhao:2010, Li:2010}.  The formation of cosmological structure in the presence of coupled scalar fields at the linear perturbative level has already been studied in great detail. Here, we extend these studies into the nonlinear regime. In particular, we want to know how structure formation on the scales of galaxies and galaxy clusters is modified. As we shall see below, there are principally four effects of the scalar field coupling, namely through (i) the modification of the background expansion rate, (ii) the action of a fifth force on dark-matter particles, (iii) the reduced contribution of dark matter to the Poisson equation, and (iv) the change of the initial condition at early times. Of these four, (ii) can actually be subdivided into (a) an essential \"rescaling\" of gravitational constant and (b) a velocity-dependent acceleration term, while (iv) is a combined consequence of (i, ii, iii) before $z\\sim 50$. We cannot track all these effects into the nonlinear regime using linear analysis and the tool we will employ to obtain quantitative predictions is $N$-body simulation \\cite{Bertschinger:1998}.  There has been some earlier work in this area (see for example,\\emph{~}\\cite% {Linder:2003, Mainini:2003, Springel:2007, Kesden:2006, Farrar:2007, Keselman:2009, Maccio:2004, Baldi:2008, Rodriguez-Meza:2007, Rodriguez-Meza:2008, Rodriguez-Meza:2009, Hellwing:2009} and also \\cite% {Oyaizu:2008a, Oyaizu:2008b, Laszlo:2008, Schmidt:2009, Chan:2009, Zhang:2010} for related works). However, in all those scalar-field simulations the scalar field equation is not solved explicitly, but rather it is assumed either that the fifth force is proportional to gravity so that its effect is a simple rescaling of the gravitational constant, or that the fifth force takes the Yukawa form with exponential cut-off beyond some scale. On the other hand, the recent fully consistent simulations performed in \\cite{Li:2009sy, Zhao:2010, Li:2010} have shown that, at least for the chameleon scalar field models, the above simplifying approximations are not good and raises questions about the extent to which we should trust them. It is these concerns that motivates this paper, in which we test the accuracy of those approximations, and study in more detail the qualitative and quantitative effects of a coupled scalar field on the formation and evolution of the nonlinear cosmic structure.  As in \\cite{Li:2010, Maccio:2004}, we shall consider a universe which at late times is dominated by a scalar field and two matter species, namely the dark matter, which couples to the scalar field in a way prescribed in Sect.~% \\ref{subsect:basiceqn} and the so-called \"baryons\", which are essentially the dark matter \\emph{without} scalar-field coupling.  The paper is organized as follows: in Sect~\\ref{sect:eqn} we list the essential equations to be implemented by our $N$-body simulations and describe briefly their differences from standard $\\Lambda $CDM. Sect.~\\ref% {sect:simu} presents a comprehensive description of the methodology used in this paper and its implementation in the numerical code. Sect.~\\ref% {subsect:code} introduces the code, Sect.~\\ref{subsect:param} lists the physical and simulation parameters we adopt, Sect.~\\ref{subsect:baryons_inc} illustrates how we distinguish between baryons and dark matter particles, Sect.~\\ref{subsect:bkgd} summarizes the results for the background expansion and linear perturbation evolution, which will be referred to subsequently from time to time, and finally Sect.~\\ref{subsect:ic} and Appendix~\\ref% {appen:Zeldovich} explain in detail how to generate initial conditions for the $N$-body simulations which take into account the effects of the coupling between dark matter and the scalar field. As we aim to set up a general framework for $N$-body simulations of coupled scalar field models, we have tried to include all the main ingredients in this section. Our numerical results are presented in Sect.~\\ref{sect:results}, within which Sect.~\\ref% {subsect:snapshots} displays some general results, showing that the approximation of rescaling gravitational constant as used in previous literature is a very good one for the model studied here (but not necessarily so for other models!), and Sects.~\\ref{subsect:Pk}, \\ref% {subsect:mf} and \\ref{subsect:halo} discuss, respectively, how the coupling modifies the nonlinear matter power spectrum, mass function and the internal density profiles of halos. We present our conclusions in \\ref{sect:con}.  All through the paper a subscript $_{\\mathrm{B}}$ ($_{\\mathrm{D}}$) denotes a corresponding quantity for baryons (dark matter), unless otherwise stated. ",
        "conclusions": "\\label{sect:con}  Couplings between any scalar field and (some of) the matter species can affect cosmology in various ways. There are two principal effects: modified source terms in the gravitational field equations (\\emph{channel I}) and direct new interactions between particles of the coupled matter species (% \\emph{channel II}).  The effects through channel II arise in an inhomogeneous universe, but not in background cosmology. In principle, there is only effect of this sort -- the fifth force, due to the exchanges of scalar quanta between matter particles. In practice, (\\emph{e.g.}, in $N$-body simulations), the scalar field $\\varphi $ is computed in the fundamental observer's frame while the fifth force must be evaluated in individual particle frames; consequently, the fifth force is split into two parts: the part from the spatial gradient of the scalar field $\\delta \\varphi $ in the fundamental observer's frame (which is independent of the particle's peculiar velocity $v$) and the part due to frame transformation which is proportional to $v$. In a homogeneous universe, there is no spatial gradient and all matter particles are comoving with the fundamental observers, so $v=0$, and therefore the channel II effects vanish. In an inhomogeneous universe both terms are nonzero in general, leading to a net force which strengthens the attraction between particles, and enhances their gravitational clustering.  The effects arising through channel I appear both in homogeneous (via the Friedmann and Raychaudhuri equations) and inhomogeneous (via the Poisson equation) universes. They arise mainly via a renormalization of the contributions to the total density from the coupled matter species and by a new contribution from the scalar field itself. \\ The exact effects depend on the specific forms of the scalar potential and coupling function. Here, we list our main findings for an inverse power-law potential Eq.~(\\ref% {eq:potential}), with an exponential coupling Eq.~(\\ref{eq:coupling_function}% ), and the model parameters given in Sect.~\\ref{subsect:param}.  \\footnotetext{% Note in Fig.~\\ref{fig:background} (lower left panel) the expansion rate of the coupled scalar field model finally catches up with that of $\\Lambda $CDM at $a=1$. This is because at late times the dark energy (scalar field) density is higher than that in $\\Lambda $CDM.}  For the background cosmology, the contribution to the Friedmann equation from dark matter is renormalized by $C(\\varphi )<1$, and $C(\\varphi )$ decreases as the coupling constant $|\\gamma |$ increases. As a result, a stronger coupling (larger $|\\gamma |$) produces lower cosmic expansion rate during the matter-dominated era (Fig.~\\ref{fig:background})\\footnotemark[5] , which will in turn helps enhance the clustering of matter and promotes structure formation.  Such an enhancement is just what we have observed from our linear perturbation analysis (Figs.~\\ref{fig:linpert} and \\ref{fig:linpert2}): on very large scales, which are beyond the range of the fifth force effects, the matter power spectrum increases with $|\\gamma |$. On small scales, the fifth force helps to enhance the matter power, making it increase further than predicted in $\\Lambda $CDM (comparing the left to the right panels of Fig.~\\ref{fig:linpert2}, where we have separated baryons, which have no scalar coupling, from dark matter, which has). This enhancement starts at a rather early, and so the scalar field does leave imprints on the matter power spectrum at $z\\sim 50$, which means that the initial condition for $N$% -body simulations also needs to modified. For example, in the model with $% | \\gamma |=0.15$, we have found that :  (i) The density perturbation at $z\\sim 49$ is about $10\\%$ larger than the corresponding result in $\\Lambda $CDM for baryons (due to slower cosmic expansion) and about $12\\%$ larger for dark matter (due to the slower expansion \\emph{and} the fifth force effects);  (ii) The average velocity at the same time is about $8\\%$ larger than the corresponding result in $\\Lambda $CDM for baryons and about $10\\%$ larger for dark matter. We introduce a quick method to take these changes into account when generating initial conditions for $N$-body simulations.  One of the most important results from our $N$-body simulations is the magnitude of fifth force. Fig.~\\ref{fig:force} shows that everywhere and at any time (\\emph{i.e.}, in both high and low density regions) the fifth force (or more precisely, the velocity-\\emph{in}dependent part of it) is proportional to gravity in magnitude (with a coefficient $2\\gamma ^{2}$ when only considering dark matter). This is the approximation used in many previous simplified simulations, and our results confirm numerically that this approximation works fairly well, at least for those cases where the scalar potential is unimportant, such as ours. We emphasize, however, that the velocity-dependent part of the fifth force should \\emph{not} be dropped (as in some previous simulations) unless there is good reason to do so (see \\cite{Li:2010}).  For the nonlinear matter power spectrum (Fig.~\\ref{fig:power}), we find that the $N$-body simulations show agreement with linear perturbation analysis on large scales. On intermediate scales, our simulations predict significant enhancement in the matter power over that predicted both by linear perturbation analysis and by the $\\Lambda $CDM paradigm. This enhancement is strongest at early epochs but gradually weakens at late times. One possible reason for this is that the scalar coupling, $C(\\varphi ),$ decreases with time, and reduces the contribution of dark matter to the gravitational potential, so weakening further clustering, but this still needs to be investigated in more detail.  The bias between dark matter and baryons increases with time and with coupling strength $|\\gamma |$. This is because larger $|\\gamma |$ implies stronger fifth forces and therefore stronger total forces act on dark matter particles compared to baryons. As time passes, the bias has more time to develop and so increases as well (Fig.~\\ref{fig:powerbc}). Using the same reasoning, the fifth force increases the attraction of dark matter particles and so heavier halos are expected to form during the structure formation (Fig.~\\ref{fig:mf}). In order to identify gravitationally bound and virialized halos in our models we must also take into account this scalar coupling (Sect.~\\ref{subsect:halo}).  The scalar field coupling also has impacts on the internal density profiles of the halos (Figs.~\\ref{fig:profile} and \\ref{fig:profiler}). We average ten of the heaviest halos and find that  (i) When compared with $\\Lambda $CDM, the overdensities in the inner regions can be lower while those in the outer regions is higher, showing the failure of the halos to retain particles in the inner region, either because the particles move too fast or because the attractive potential at the centre is too weak;  (ii) The suppression of inner density compared with $\\Lambda $CDM weakens, rather than strengthens, as $|\\gamma |$ increases. Indeed, for $|\\gamma | =0.2 $ the density is higher than $\\Lambda $CDM result throughout the halos. More detailed studies are needed to clarify the leading effect responsible for this observed pattern.  (iii) There is a bias between density profiles for baryons and dark matter (Fig.~\\ref{fig:profilebc}): the dark matter density is always higher because it experiences the fifth force which boosts its clustering.  In summary, in this paper we have been given a comprehensive description of the methodology and implementation of general $N$-body simulations for coupled scalar field models. Some important issues in these simulations have been addressed here for the first time: the consistent solution of the scalar field and fifth force, the fifth force effects on generating initial conditions, and the effects of scalar field on identifying virialized halos. Although the situation is complex, we have identified interesting new features in these models. We hope these developments will lead to more detailed study of nonlinear structure formation in this class of models and facilitate their subsequent confrontation with observational data."
    },
    "1005/1005.3885_arXiv.txt": {
        "abstract": "We present a CCD photometric study of the star with ASAS ID 134738$+$0410.1 using V band observations obtained from the $IUCAA$ Girawali Observatory (IGO) 2-metre telescope, India.~The star was selected from the $\\delta$ Scuti database of All Sky Automated Survey (ASAS) (Pojmanski 2002). Our analysis reveals that the star is not a $\\delta$ Scuti variable but is in fact a W UMa type contact binary with an orbital period of 0.2853067 day. Two new times of primary and secondary minima were determined from the observed data. A preliminary solution obtained using the Wilson-Devinney light curve modelling technique indicates that the star is more likely a partially-eclipsing W UMa type contact binary. However, the determination of actual subtype of this binary is quite impossible from the photometry alone, as the observed light curve can fitted for both A- and W-type solutions. The exact classification of this binary needs to be determined from high resolution spectroscopy. ",
        "introduction": "\\label{sec:intro} In the ASAS\\footnote{http://www.astrouw.edu.pl/asas/} database, the star has maximum V magnitude of 13.47 mag and is located at $\\alpha = 13^{\\rm h}\\, 47^{\\rm m}\\,38^{\\rm s}.0$ and $\\delta = 04^{\\rm o}\\,10^{\\rm '}\\,05^{\\rm ''}.9$ (Epoch J2000). The star has been classified as $\\delta$ Scuti/Contact/ Semi-Detached binary with a period of P$_{\\rm ASAS}$ = 0.142654 day. The classification is thus ambiguous. In order to know true nature of this object, we have obtained high precision CCD photometric data from new observations. The new data indicates that the star is a W UMa contact binary with an orbital period of P$_{\\rm IGO} = 0.2853067$ day which is nearly twice the period P$_{\\rm ASAS}$ = 0.142654 day.  \\begin{table*} \\begin{center} \\caption{Basic information of ASAS 134738$+$0410.1, the comparison star and the check star} \\label{Table1} \\scalebox{0.82}{ \\begin{tabular}{lccccc} \\hline Star  & RA (J2000) & DEC (J2000) & J [mag] & H [mag]& K [mag]  \\\\ \\hline ASAS 134738$+$0410.1 (Target) & $13^{\\rm h}\\,47^{\\rm m}\\,38^{\\rm s}.00$ & $04^{\\rm o}\\,10^{\\rm '}\\, 06^{\\rm ''}.00$ &11.894 &11.472 &11.407  \\\\  2MASS 13473475$+$0408167 (Comparison) & $13^{\\rm h}\\,47^{\\rm m}\\,34^{\\rm s}.76$ & $04^{\\rm o}\\,08^{\\rm '} 16^{\\rm ''}.72$ &12.408&12.026 &11.951\\\\ 2MASS 13472306$+$0413420 (Check) & $13^{\\rm h}\\,47^{\\rm m}\\,23^{\\rm s}.06$ & $04^{\\rm o}\\,13^{\\rm '}\\,42^{\\rm ''}.10$ &12.213 &11.870 &11.789 \\\\ \\hline \\end{tabular}}  \\end{center} \\end{table*}  The V band CCD photometric observations of the star were carried out with the IGO 2-m telescope, located about 80 km from Pune, India during two nights on March 31 and April $02$  in 2009.~The IUCAA Faint Object Spectrograph Camera (IFOSC) equipped with EEV 2 K$\\times$ 2 K thinned, back-illuminated CCD with 13.5 $\\mu m$ pixels was used. The CCD used for imaging provides an effective field of view of $\\sim 10.5^{'} \\times 10.5^{'} $ on the sky corresponding to a plate scale of 0.3 arcsec pixel$^{-1}$. The gain and read out noise of the CCD camera are 1.5 e$^{-}$/ADU and 4 e$^{-}$ respectively. The FWHM of the stellar image varied from 3 to 5 pixels during the observations. We took a total of 153 frames in the V band with the exposure times varied between 100 s and 180 s for a good photometric accuracy.  The co-ordinates of the variable, comparison star and the check star along with the infrared JHK magnitudes taken from the 2MASS catalogue (Cutri et al. 2003) are listed in Table 1. The comparison and the check star are so close to the variable that they are in the same field during the observations. Image pre-processing and data reduction was carried out using IRAF\\footnote {IRAF is distributed by the National Optical Astronomy Observatories, which are operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.} and MIDAS softwares. Instrumental magnitudes were obtained using the DAOPHOT package (Stetson 1987, 1992). The various tasks, e.g., $find, phot, daogrow, daomatch$ and $daomaster$ were applied in order to obtain the instrumental magnitudes of stars in all the frames. Extinction corrections were ignored as the target star is very close to the comparison star. In Fig.~1, we show the plots of the differential V band magnitude of (Variable - Comparison), (Comparison - Check) versus Heliocentric Julian Day (HJD) in the left, right upper and lower panels respectively for observations on March 31 and April 02, 2009. The reduced results show that the difference between the magnitude of the check star and that of the comparison star was constant with a probable error of $\\pm$0.004 mag in the V band \\footnote{The observational data presented in this paper can be obtained from the authors on request.}. ",
        "conclusions": "We have discovered that star having ASAS ID 134738+0410.1 identified by ASAS is a W UMa type overcontact binary.  We have determined new period and times of minima from very accurate and precise CCD data. Based on the WD code as implemented in the software PHOEBE, we have done preliminary modelling of the light curve. However, unique value of the mass ratio could not be obtained from the modelling, due to its partial-eclipsing nature. Hence, the actual subtype of the system could not be determined. This star deserves high resolution spectroscopic time series radial velocity measurements to determine its true subtype and mass ratio which should be combined with the photometric light curve data to obtain various geometrical and physical parameters accurately."
    },
    "1005/1005.5147_arXiv.txt": {
        "abstract": "{We present results from the earliest observations of DEBRIS, a {\\it Herschel} Key Programme to conduct a volume- and flux-limited survey for debris discs in A-type through M-type stars.  PACS images (from chop/nod or scan-mode observations) at 100 and 160 \\micron\\ are presented toward two A-type stars and one F-type star: $\\beta$ Leo, $\\beta$ UMa and $\\eta$ Corvi. All three stars are known disc hosts.  {\\it Herschel} spatially resolves the dust emission around all three stars (marginally, in the case of $\\beta$ UMa), providing new information about discs as close as 11 pc with sizes comparable to that of the Solar System.  We have combined these data with existing flux density measurements of the discs to refine the SEDs and derive estimates of the fractional luminosities, temperatures and radii of the discs.  } ",
        "introduction": "\\label{intro}  Debris discs are flattened distributions of planetesimals and dust located at radii of 1-1000~AU around main-sequence stars (see \\citep[see][for a recent review]{wyatt08}. The dust cannot be primordial since its lifetime in orbit is significantly less than the age of the host stars.  Instead, dust is replenished from a population of colliding km-sized planetesimals \\citep{wyatt02,the07}.  Over time, the dust distribution is shaped by any planetary-sized bodies in the system \\citep[e.g.,][]{dom03,wyatt07}. Therefore, resolved images of discs constrain models of the structure and evolution of planetary systems.  Far-infrared and submillimetre observations are the best way to search for dust around nearby stars due to the favorable contrast of the disc relative to the star.  At these wavelengths, the disc emission is optically thin and is sensitive to the large (up to $\\sim 1$ mm), cold grains which dominate the disc's dust mass.  The {\\em Herschel} Space Observatory offers three major advantages for the detection and characterization of debris discs: far-infrared sensitivity, angular resolution and wavelength coverage. With its 3.5 m mirror, its sensitivity at far-infrared wavelengths is superior to any previous instrument. With its resolution of 6\\farcs7 at 100 \\micron, {\\it Herschel} has the potential to resolve many debris discs, particularly toward nearby stars. Finally, with detectors at 100, 160, 250, 350 and 500 \\micron, {\\it Herschel} has the means to sample the spectral energy distribution (SED) of disc emission across the peak, meaning models can be better constrained even for discs which are not resolved.  DEBRIS (Disc Emission via a Bias-free Reconnaissance in the Infrared/Submillimetre) is an Open Time Key Programme which uses PACS (Photodetector Array Camera \\& Spectrometer) and (for appropriate targets) SPIRE (Spectral and Photometric Imaging Receiver) to detect, resolve and characterize debris discs around a volume-limited sample of 446 A through M type stars.  The goals of DEBRIS include establishing the incidence and evolution of debris discs as a function of stellar type, age, multiplicity, etc.; the characterization of discs in terms of size, temperature, dust mass and morphology (where the disc asymmetries could indicate the presence of planetary companions); and the understanding of our own Solar System in the context of the larger debris disc population.  Full details of the DEBRIS survey and goals will be presented in a forthcoming paper (B.~Matthews et al. 2010, in preparation).  Here, we present PACS observations toward three of the first targets of the DEBRIS survey. We briefly summarize the observations and targets in Sect.~\\ref{obs}, present the results in Sect.~\\ref{res} and discuss three sources in detail in Sect.~\\ref{disc}. We summarize the paper in Sect.~\\ref{sum}.    \\begin{figure*} \\centering \\includegraphics[bb=20 335 640 645,width=20cm,clip]{matthewsb_fg1.ps} \\caption{Images of the 100 and 160 \\micron\\ emission from three DEBRIS targets: $\\beta$ Leo, $\\beta$ UMa and $\\eta$ Corvi. Contours are shown at 0, 10, 30, 50, 60, 70, 80, 90 and 99\\% of the peak in each map. The 1-$\\sigma$ rms noise levels are given in Table \\ref{table}. Circles in the upper left corner of each panel mark the nominal beam sizes for a scan speed of 20\\arcsec/s, i.e., 6\\farcs7 and 11\\arcsec\\ at 100 and 160 \\micron, respectively. The negative images created by the chop/nod observing mode are visible in the $\\beta$ UMa images. Striping in the $\\eta$ Corvi image at 160 \\micron\\ is due to high 1/f noise and filtering artefacts.  } \\label{images} \\end{figure*} ",
        "conclusions": "\\label{disc}  \\subsection{$\\beta$ Leo} \\label{betaleo}   Figure \\ref{images} shows the first resolved images of the disc around $\\beta$ Leo.  The blackbody temperature and dust luminosity results given in Table \\ref{table} are consistent with the values found in previous works \\citep{su06,holmes03}.  The fractional dust luminosity of $2.3 \\times 10^{-5}$ is 15\\% higher than the estimate from \\cite{su06}.  The increase is due to a slight increase in dust to match the PACS flux densities.  The radius estimates found for $\\beta$ Leo are comparable to that of the Kuiper Belt ($\\sim 50$ AU). This makes the $\\beta$ Leo disc one of the smallest disc radii yet resolved at any wavelength (see, for instance, the 'Circumstellar Disks Database'\\footnote{cirumstellardisks.org}) although smaller characteristic orbital radii have been derived based on single temperature blackbody fits to the dust components \\citep[e.g.,][]{rhee07}.  The difference between $R_{obs}$ and $R_{dust}$ provides an opportunity to learn about the grains within this disc.  Because differently sized grains can have the same temperature at different distances from a star, the SED models in Fig.\\ \\ref{SEDs} are degenerate. This degeneracy is broken by the resolved imaging. For example, the $\\sim$40 AU radius for the $\\beta$ Leo disc is larger than the 23 AU suggested by the blackbody fit. Therefore, the grains do not emit as blackbodies, but maintain a $\\sim$112 K temperature at a greater distance from the star as expected for small grains that emit inefficiently at far-IR wavelengths.  The inferred characteristic particle radius $a$ is $ < \\lambda/2 \\pi = 16$ \\micron. Future modeling work that combines {\\it Spitzer} IRS spectra and submillimetre images with the {\\it Herschel} data will constrain these grain properties and the spatial dust distribution (L.~Churcher et al. 2010, in preparation).  \\subsection{$\\beta$ UMa} \\label{betauma}  Figure \\ref{images} shows that the disc emission around $\\beta$ UMa is very compact at 100 and 160 \\micron.  The disc is marginally resolved at 100 \\micron\\ and not resolved at 160 \\micron. The apparent asymmetry in the 160 \\micron\\ disc image is likely artificial; it is an effect of interpolation applied to the image at native (Nyquist sampled) resolution.  The flux densities measured for $\\beta$ UMa confirm the earlier 100 and 160 \\micron\\ detections.  The disc component of the $\\beta$ UMa SED is well fit by blackbody grains with a temperature comparable to that of $\\beta$ Leo, requiring a bigger disc around the more luminous star.  Therefore, assuming black body grains the radial estimate is 51AU, equivalent to the deconvolved disc radius from the 2D Gaussian fit to the 100 \\micron\\ image.  The resolved size thus suggests an absence of small grains such as that inferred for beta Leo in Sect.~\\ref{betaleo}. More detailed modeling will be forthcoming in a future paper.  \\subsection{$\\eta$ Corvi}  The new {\\it Herschel} images in Fig.\\ \\ref{images} show that $\\eta$ Corvi is resolved at both 100 and 160 \\micron, as expected based on the $\\sim 300$ AU submillimeter size derived by \\cite{wyatt05}.  The variation in morphology from centrally peaked emission at 100 \\micron\\ to a double-peaked limb brightened ring at 160 \\micron\\ (as observed at 450 \\micron) is consistent with an outer cool ring filled in by warmer dust which dominates the emission at 100 \\micron. This could be evidence of the third temperature component proposed by \\cite{chen06} and observed in $\\epsilon$ Eri by \\cite{back09}, although this was tentatively ruled out in mid-IR imaging by \\cite{smith08}, and more generally suggests the radial distribution of material is broader than the two ring system originally envisaged by \\cite{wyatt05}.  The $R_{obs}$ estimate from Table \\ref{table} is equivalent to the submillimetre size. The two intensity maxima in the 160 \\micron\\ image are roughly a beamwidth (11\\arcsec) apart, identical to the 450 \\micron\\ SCUBA imaging of \\cite{wyatt05} who inferred that the emission arises from a ring at moderate inclination. Fitting a 2D Gaussian to the 100 \\micron\\ image of Fig.\\ \\ref{images} gives a position angle of 102\\degr $\\pm 7$\\degr and an inclination of $\\sim$ 50\\degr\\ from the line of sight. The position angle of the two peaks at 160 \\micron\\ is $\\sim 135$\\degr. These inclinations are consistent with the 450 \\micron\\ measurement of $130$\\degr\\ $\\pm 10$\\degr.  A two component model of the SED of Fig.\\ \\ref{SEDs} shows similar results to \\cite{wyatt05} who found disc components of 40 K and 370 K. The warm component \\citep{smith09,smith08,chen06} shown in Fig.\\ \\ref{SEDs} has a blackbody temperature of 346 K, corresponding to a radial distance of 1.4 AU from the star.  The cold component has a temperature of 33 K, corresponding to a radial separation of 160 AU from the star, consistent with $R_{obs}$.  As for $\\beta$ UMa, the resolved size suggests an absence of the small grains implied for $\\beta$ Leo in Sect.~\\ref{betaleo}.  Most importantly, the images in Fig.\\ \\ref{images} provide an estimate of the disk size at wavelengths intermediate between the submillimetre (which shows emission at $\\sim 150$ AU) and mid-IR (which shows emission at $<$ 3.5 AU). This will be crucial for modeling the origin of the far-infrared morphology, which most resembles the submillimetre emission.  Simultaneously modeling these several images will constrain in more detail the dust properties of the disc system."
    },
    "1005/1005.2365.txt": {
        "abstract": "In the past few years, several disks with inner holes that are relatively empty of small dust grains have been detected and are known as transitional disks.  Recently, {\\it Spitzer} has identified a new class of ``pre-transitional disks\" with gaps based on near-infrared photometry and mid-infrared spectra; these objects have an optically thick inner disk separated from an optically thick outer disk by an optically thin disk gap. A near-infrared spectrum provided the first confirmation of a gap in the pre-transitional disk of LkCa~15 by verifying that the near-infrared excess emission in this object was due to an optically thick inner disk. Here we investigate the difference between the nature of the inner regions of transitional and pre-transitional disks using the same veiling-based technique to extract the near-infrared excess emission above the stellar photosphere. However, in this work we use detailed disk models to fit the excess continua as opposed to the simple blackbody fits used previously. We show that the near-infrared excess emission of the previously identified pre-transitional disks of LkCa~15 and UX~Tau~A in the Taurus cloud as well as the newly identified pre-transitional disk of Rox~44 in Ophiuchus can be fit with an inner disk wall located at the dust destruction radius. We also present detailed modeling of the broad-band spectral energy distributions of these objects, taking into account the effect of shadowing by the inner disk on the outer disk, but considering the finite size of the star, unlike other recent treatments. The near-infrared excess continua of these three pre-transitional disks, which can be explained by optically thick inner disks, are significantly different from that of the transitional disks of GM~Aur, whose near-infrared excess continuum can be reproduced by emission from sub-micron-sized optically thin dust, and DM~Tau, whose near-infrared spectrum is consistent with a disk hole that is relatively free of small dust. The structure of pre-transitional disks may be a sign of young planets forming in these disks and future studies of pre-transitional disks will provide constraints to aid in theoretical modeling of planet formation. ",
        "introduction": "\\label{intro}  Several disks which have nearly photospheric near-infrared emission but substantial excesses above the stellar photosphere at wavelengths beyond $\\sim$20~{\\micron} have been observed and are referred to as ``transitional disks'' \\citep{strom89}. Using data from the {\\it Spitzer} Infrared Spectrograph \\citep[IRS;][]{houck04}, detailed modeling has demonstrated that this flux deficit at near-infrared wavelengths relative to full disks can be explained by optically thick disks with inner holes of $< 40$~AU. In most cases these inner holes are not completely devoid of material (e.g. GM~Aur, TW Hya, CS Cha, CVSO 224); a minute amount of micron- or sub-micron-sized optically thin dust exists within the hole, producing a small infrared excess over the photospheric flux, still well below the median excess of Class II objects, as well as silicate emission \\citep{calvet02,calvet05, espaillat07a, espaillat08b}.  Gas has also been detected within the inner holes of transitional disks \\citep[e.g.][]{najita03,bergin04,salyk07}.  Recently, the {\\it Spitzer Space Telescope} \\citep{werner04} identified a new class of disks called ``pre-transitional disks\" around LkCa~15 and UX~Tau~A \\citep{espaillat07b}. These disks have deficits of mid-infrared flux (5--20~{\\micron}) and substantial excesses at longer wavelengths, as is seen in the transitional disks. However, in contrast to the small or absent near-infrared (2--5~{\\micron}) excesses exhibited by transitional disks, pre-transitional disks have significant near-infrared excesses relative to their stellar photospheres, similar to the median spectral energy distribution (SED) of disks in Taurus \\citep{dalessio99}.  The distinctive shapes of these SEDs indicate that pre-transitional disks have an inner disk separated from an outer disk and that we may be seeing the development of gaps within protoplanetary disks.  While the truncation of LkCa~15's outer disk has been imaged in the millimeter \\citep{pietu06}, \\citet{espaillat07b} showed that the substantial near-infrared excess of LkCa~15 could be explained by either optically thick material or by ${\\sim}10^{-11}$ M$_{\\odot}$ of optically thin dust mixed with the gas in the inner disk.  In order to resolve this issue, \\citet{espaillat08a} obtained a medium resolution near-infrared spectrum spanning the wavelength range 2--5~{\\micron}. This near-infrared spectrum had absorption lines that were weaker relative to the spectrum of a standard star of the same spectral type. This phenomenon, known as ``veiling\" \\citep{hartigan89}, is also observed in similar spectra of full disks and is due to emission from dust located at the dust sublimation radius \\citep{muzerolle03}. \\citet{espaillat08a} measured a veiling factor (r$_{K}$) of 0.3$\\pm$0.2 for LkCa~15 at $\\sim$2.2~{\\micron} and fit the near-infrared excess from 2--5~{\\micron} with a single-temperature blackbody of 1600~K.  This behavior can be explained by an optically thick wall located at the dust sublimation radius, as is seen in full disks \\citep{muzerolle03}. These data confirmed that LkCa~15 has an inner optically thick disk, making this observation the first independent verification of a gap in a protoplanetary disk.  Here we expand our sample to include the pre-transitional disks of UX Tau A \\citep{espaillat07b} and Rox~44 \\citep{furlan09},  and the transitional disks of GM~Aur and DM~Tau \\citep{calvet05}  in order to explore the structure of the inner regions of pre-transitional and transitional disks.  To do this we use veiling measurements in the K-band to extract the near-infrared (-IR) excess emission of these objects and then fit this emission with disk models \\citep{dalessio05,calvet02}. Given the sensitivity of veiling measurements on the adopted spectral type, we redetermined spectral types for our targets using optical spectra ({\\S}~\\ref{sec:spt}). We also tested our veiling measurement methods with the spectrum of the diskless, weak-line T Tauri star (WTTS) LkCa~14 (\\S~\\ref{sec:lkca14}).  We find that the near-infrared spectra of the pre-transitional disks of LkCa~15, UX~Tau~A, and Rox~44 are well-explained by the wall of an optically thick inner disk ({\\S}~\\ref{sec:lkca14}, {\\S}~\\ref{sec:ux}, {\\S}~\\ref{sec:rox}). In contrast, our data shows that the inner hole of the transitional disk of GM~Aur contains a small amount of optically thin sub-micron-sized dust while DM~Tau's hole is relatively free of small dust ({\\S}~\\ref{sec:gm} \\& {\\S}~\\ref{sec:dm}). Our results are consistent with veiling and interferometric measurements found in the literature. We also perform detailed model fits to the broad-band SEDs of our pre-transitional disk sample, and explore the effect of shadowing of the outer disk by the inner disk ({\\S}~\\ref{sec:models}), taking into account that the star is not a point source, as has been adopted in other studies \\citep{espaillat07b, mulders10}. The structure of pre-transitional disks suggests that of the disk clearing mechanisms proposed to date, planet formation \\citep[e.g.][]{goldreich80, rice03, varniere06} is most likely a dominant factor in clearing these disks. ",
        "conclusions": "\\label{sed:sum}   In this paper we presented spectral type measurements as well as veiling and near-infrared excess measurements for the diskless LkCa~14, the pre-transitional disks of LkCa~15, UX~Tau~A, and Rox~44, and the transitional disks of DM~Tau and GM~Aur. Using near-IR SpeX spectra from 1--5 {\\micron}, we extracted the near-IR excess continua of our pre-transitional disk sample and fit their excess emission with optically thick inner disk walls, supporting the interpretation that these objects contain gaps in their disks. We modeled the broad-band SEDs of these pre-transitional disks and measure gap sizes of $\\sim$40--70~AU.  In the case of LkCa~15 and Rox~44 there is $\\sim$10$^{-11}$~{\\msun} of ISM-sized optically thin dust within the gap.  The near-infrared emission of our pre-transitional disks differs significantly from that of the transitional disks of DM~Tau and GM~Aur. DM~Tau has no veiling and no excess emission in the near-IR, indicating that its disk hole is relatively empty of small dust grains.  GM~Aur has a small amount of veiling and displays a flat excess above the stellar photosphere in the near-IR, which can be fit by a model of optically thin dust emission.  This is in contrast to the pre-transitional disks which have large, blackbody-like near-IR excess continua that can be fit with models of a wall located at the dust sublimation radius.  We also studied the effects of shadowing of the outer disk wall by the inner disk. We found that when the finite size of the star is taken into account, a significant portion of the outer wall is either in the prenumbra or completely out of the shadow of the inner wall. The predicted height of the wall is consistent with that of the outer disk, which is derived independently.  Based on currently known disk clearing mechanisms, we propose that the gaps in pre-transitional disks are indicators of planet formation, making these disks a promising location for young planet searches. {\\it ALMA} will be pivotal in extending the sample of known pre-transitional disks and has the potential to detect planets in these gaps. Near-IR interferometry will play an important role in further understanding the innermost regions of these disks. Future studies of this class of objects may bring us a few steps closer to understanding the origin of our own solar system."
    },
    "1005/1005.5001_arXiv.txt": {
        "abstract": "s{ Excitation of multicomponent dark matter in the galactic center has been proposed as the source of low-energy positrons that produce the excess 511 keV $\\gamma$ rays that have been observed by INTEGRAL.  Such models have also been promoted to explain excess high-energy $e^\\pm$ observed by the PAMELA, Fermi/LAT and H.E.S.S.\\ experiments.  We investigate whether one model can simultaneously fit all three anomalies, in addition to further constraints from inverse Compton scattering by the high-energy leptons.  We find models that fit both the 511 keV and PAMELA excesses at dark matter masses $M < 400$ GeV, but not the Fermi lepton excess.  The conflict arises because a more cuspy DM halo profile is needed to match the observed 511 keV signal than is compatible with inverse Compton constraints at larger DM masses. } ",
        "introduction": " ",
        "conclusions": "We have found that annihilating multistate DM can explain two out of three galactic cosmic ray anomalies, either PAMELA/Fermi or PAMELA/INTEGRAL, but not all three simultaneously.   Although it is possible to marginally predict all the correct rates using Einasto profile parameter $\\alpha=0.20$, the angular distribution of 511 keV $\\gamma$ rays is too wide in this case.   Of the two possibilities, the PAMELA/INTEGRAL combination seems preferable from the standpoint of the required DM halo parameters, since in this case we are able to adopt standard values that are quite compatible with $N$-body simulations of galactic structure evolution.  Moreover we can match the anomalous lepton rates well for PAMELA/INTEGRAL.  The PAMELA/Fermi possibility requires stretching the halo parameters to their maximal values, while only marginally giving a large enough rate of leptons, yet a small enough rate of associated inverse Compton $\\gamma$ rays."
    },
    "1005/1005.2621_arXiv.txt": {
        "abstract": "We present the discovery of an extremely bright and extended lensed source from the second Red Sequence Cluster Survey (RCS2). RCSGA 032727-132609 is spectroscopically confirmed as a giant arc and counter-image of a background galaxy at $z=1.701$, strongly-lensed by the foreground galaxy cluster RCS2 032727-132623 at $z=0.564$. The giant arc extends over $\\sim 38$\\,\\arcsec and has an integrated $r$-band magnitude of 19.1, making it $\\sim 20$ times larger and $\\sim 3.5$ times brighter than the prototypical lensed galaxy MS1512-cB58. This is the brightest distant lensed galaxy in the Universe known to date. % We have collected photometry in 9 bands, ranging from $u$ to $K_s$, which densely sample the rest-frame UV and optical light, including the age-sensitive 4000\\AA\\ break. A lens model is constructed for the system and results in a robust total magnification of $2.04 \\pm 0.16$ for the counter-image; we estimate an average magnification of $17.2 \\pm 1.4$ for the giant arc based on the relative physical scales of the arc and counter-image on the sky. Fits of single-component spectral energy distribution (SED) models to the photometry result in a moderately young age, $t=80\\pm40$\\,Myr, small amounts of dust, $E(B-V) \\le 0.11$, and an exponentially declining star formation history with \\textit{e}-folding time $\\tau = 10-50$\\,Myr. After correcting for the lensing magnification, we find a stellar mass of M$_* \\sim 10^{10}$\\,M$_\\odot$ and a current star formation rate SFR$\\le77$\\,M$_\\odot$\\,yr$^{-1}$. Allowing for episodic star formation, an underlying old burst could contain up to twice the mass inferred from single-component modeling. % RCSGA 032727-132609 is typical of the known population of star-forming galaxies near this redshift in terms of its age and stellar mass. Its large magnification and spatial extent provide a unique opportunity to study the physical properties of an individual high-redshift star-forming galaxy in great detail, opening up a new window to the process of galaxy evolution between $z=1.7$ and our local Universe.  \\subjectheadings{galaxies: evolution, galaxies: high-redshift, gravitational lensing} ",
        "introduction": "Significant progress has been made in recent years towards the study of the formation of galaxies and their evolution into the population of objects we observe around us today. The redshift range $1.0 \\lesssim z \\lesssim 3.0$ represents a crucial period in this process since the peak of star formation in the Universe is known to occur near $z \\sim 2$ \\citep{Blain:99, Chapman:03, Reddy:09}. Through the development of pre-selection color criteria \\citep{Adelberger:04, Steidel:04, Daddi:04a} and advances in the near-IR and UV spectroscopic capabilities of 8-10\\,m class telescopes, growing samples of optically or near-IR selected galaxies are now spectroscopically confirmed in this redshift range and studied extensively (see e.g. Shapley et al. 2005; Erb et al. 2006; Kriek et al. 2008). For a complete understanding of the process of galaxy formation, it is crucial to complement these statistical results based on larger samples with detailed study of the stellar populations and dynamics of individual objects. The majority of galaxies in the current samples are too faint for this purpose. The few galaxies that are bright enough to be studied individually represent outliers drawn from the extreme bright tail of the luminosity function, and are therefore not necessarily representative of the bulk of the population. The ability to study the properties of faint, high-redshift galaxies is one of the main science drivers for the construction of 30\\,m class telescopes. Until such instruments become available we can make a head start using gravitational lensing to increase the power of the current generation of telescopes.  The lensing magnification induced by foreground galaxy clusters and individual galaxies has been succesfully used to identify galaxies out to $z \\sim 10$, opening up new windows into the very distant Universe \\citep{Richard:08, Bouwens:09}. At more moderate redshifts it brings individual galaxies from the samples at $1.0 \\lesssim z \\lesssim 3.0$ to a flux level amenable for extensive follow-up observations at various wavelengths. The first and most notable example in this class is MS1512-cB58, a Lyman Break Galaxy (LBG) at $z = 2.73$ \\citep{Yee:96, Ellingson:96}, found to have a lensing magnification of $\\sim$ 30 \\citep{Williams:96, Seitz:98}. It has been studied extensively since its discovery and provides a wealth of information on the stellar population and dynamics of a young star-forming galaxy at this redshift \\citep{Pettini:00, Pettini:02, Siana:08}. Other examples of particularly bright lensed galaxies at similar redshifts include the `Cosmic Eye' at $z=3.07$ \\citep{Smail:07,Coppin:07}, the `8 o'clock arc' at $z=2.73$ \\citep{Allam:07, Finkelstein:09} and two strongly-lensed $z\\sim3$ LBGs from the SDSS Giant Arcs Survey \\citep{Koester:10}. These highly magnified sources currently represent the best places to study the individual properties of high-redshift star-forming galaxies in great detail. %  In this paper we present the discovery of a very bright and extended galaxy, RCSGA 032727-132609, spectroscopically confirmed at $z = 1.701$ and highly magnified by a foreground cluster at $z = 0.564$ from the second Red Sequence Cluster Survey (RCS2; D. Gilbank et al. 2010, in preparation). The RCS surveys were designed for the purpose of cluster finding via the identification of the linear color-magnitude relation present for early-type galaxies in clusters, known as the red-sequence technique \\citep{Gladdersyee:00}. The RCS2 survey has imaged $\\sim700$ square degrees of sky in $g$,$r$ and $z$ with the MegaCam camera at the Canada-France-Hawaii Telescope (CFHT) on Mauna Kea in 4, 8 and 6\\,min exposures respectively. Data acquisition for the survey finished in 2008 and preliminary cluster catalogs over the entire area have been created and visually inspected for strong lensing signatures; the details of this search will be published elsewhere. The brightest and most obvious strong lensing system found in this search is RCSGA (RCS Giant Arc) 032727-132609. The system consists of a counter-image and a giant arc extending over $\\sim 38$\\,\\arcsec at an Einstein-radius of $\\sim 17.8$\\,\\arcsec, estimated from the distance between the brightest knot in the arc and the brightest cluster galaxy. The arc has an apparent magnitude of $r = 19.1$, making it $\\sim 3.5$ times brighter than cB58 and the brightest distant lensed galaxy in the Universe known to date. Its large spatial extent provides unique opportunities to look inside a high-redshift galaxy and spatially resolve its substructure. \\npar The paper is organized as follows. \\S~\\ref{sec:data} presents the multitude of photometric and spectroscopic data we have assembled on this system: broadband observations in 9 bands ranging from $u$ to $K_s$, a medium-resolution optical spectrum of the lensed galaxy to obtain the source redshift and redshift measurements of 49 cluster members to estimate the virial mass of the foreground cluster. \\S~\\ref{sec:phot} describes an innovative method used to obtain accurate photometry of the source. A lens model is constructed for the cluster in \\S~\\ref{sec:lensmodel}. We fit spectral energy distribution (SED) models to the broadband photometry of the source to explore its star formation history and stellar population parameters; the SED modeling procedure and results are discussed in \\S~\\ref{sec:sedall}. RCSGA 032727-132609 is compared to the known galaxy population at $z\\sim 2$ in \\S~\\ref{sec:compall}. Throughout this paper we will assume $\\Omega_M = 0.3$, $\\Omega_\\Lambda = 0.7$ and H$_0 = 70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. All magnitudes are quoted in the AB system. ",
        "conclusions": "This paper reports on the discovery of an exceptionally bright and extended star-forming galaxy at $z=1.701$, strongly-lensed by a foreground cluster discovered in the RCS2 survey. The giant arc is $\\sim 3.5$ times brighter than cB58 and extends over 38\\,\\arcsec on the sky. We measure a velocity dispersion of 988$\\pm$122\\,km~s$^{-1}$ for the cluster and estimate a virial mass of $M_{200} \\sim 1.1 \\times 10^{15}$\\,M$_\\odot$\\,$h_{70}^{-1}$, in accord with the large Einstein-radius of $\\sim 17\\farcs8$. A lens model is constructed for the cluster using the publicly-available software {\\tt LENSTOOL} \\citep{Jullo:07} and results in a magnification factor of 2.04$\\pm$0.16 for the counter-image and an estimate of 17.2$\\pm$1.4 for the average magnification of the giant arc, based on the relative sizes of the arc and the counter-image on the sky. Higher resolution imaging is required to create a robust lens model for the giant arc which correctly explains the apparent substructure.  Careful measurements of consistent magnitudes for 9 bands of photometry ranging from $u$ to $K_s$ produce a well-constrained spectral energy distribution. SED fitting is based on CB07 models with a Chabrier IMF and Calzetti dust extinction. The systematic uncertainties which arise from the choice of stellar evolution model, IMF and extinction law can be large and often exceed the statistical uncertainties from photometric errors. Our best-fit model places RCSGA 032727-132609 at a metallicity of 0.4\\,Z$_\\odot$ with a moderately young age, $t=80\\pm40$\\,Myr and a relatively small amount of dust extinction, $E(B-V) \\le 0.11$. Taking into account the lensing magnification, we find a stellar mass of M$_* \\sim 10^{10}$\\,M$_\\odot$ and a current star formation rate $SFR \\le 77$\\,M$_\\odot$\\,yr$^{-1}$. The agreement between the stellar population parameters independently inferred from the giant arc and counter-image, both images of the same background galaxy, is a good consistency check of the photometry, SED fitting procedure and lens modeling. Single-component SED models report ages and stellar mass measurements for the current episode of star formation. When allowing for an episodic star formation history, twice this mass could be present in an older underlying burst without affecting the SED detectably.  A comparison to the characteristic L$_*$ and M$_{\\mathrm{star}}^*$ of the Schechter luminosity and mass functions at $z\\sim1.7$ places RCSGA 032727-132609 at a luminosity of $6.5\\pm1.8$\\,L$_*$ and a stellar mass of $0.16\\pm0.04$\\,M$_{\\mathrm{star}}^*$, suggesting a low-mass starbursting galaxy. The Salpeter stellar mass estimate of RCSGA 032727-132609 is consistent with the average stellar mass for a sample of `BM' galaxies at $z = 1.72 \\pm 0.34$ \\citep{Reddy:06a} and its age, reddening and current star formation rate generally fall into the broad range of stellar population parameters found for this comparison sample. The large magnification and spatial extent of RCSGA 032727-132609 provide a unique opportunity to spatially resolve one representative example of the variety of stellar populations found at these redshifts and study it in great detail. This paper has presented only our first level of understanding of this unique galaxy."
    },
    "1005/1005.0630_arXiv.txt": {
        "abstract": "% We present deep {\\it Spitzer} mid-infrared spectroscopy, along with 16, 24, 70, and 850\\,$\\micron$\\ photometry, for 22 galaxies located in the Great Observatories Origins Deep Survey-North (GOODS-N) field. The sample spans a redshift range of $0.6\\la z \\la 2.6$, 24~$\\mu$m flux densities between $\\sim$0.2$-$1.2 mJy, and consists of submillimeter galaxies (SMGs), X-ray or optically selected active galactic nuclei (AGN), and optically faint ($z_{AB}>25$\\,mag) sources. We find that infrared (IR; $8-1000~\\micron$) luminosities derived by fitting local spectral energy distributions (SEDs) with 24~$\\micron$ photometry alone are well matched to those when additional mid-infrared spectroscopic and longer wavelength photometric data is used for galaxies having $z\\la1.4$ and 24~$\\micron$-derived IR luminosities typically $\\la 3\\times 10^{12}~L_{\\sun}$. However, for galaxies in the redshift range between $1.4\\la z \\la 2.6$, typically having 24~$\\micron$-derived IR luminosities $\\ga 3\\times 10^{12}~L_{\\sun}$, IR luminosities are overestimated by an average factor of $\\sim$5 when SED fitting with 24~$\\micron$ photometry alone. This result arises partly due to the fact that high redshift galaxies exhibit aromatic feature equivalent widths that are large compared to local galaxies of similar luminosities. Through a spectral decomposition of mid-infrared spectroscopic data, we are able to isolate the fraction of IR luminosity arising from an AGN as opposed to star formation activity. This fraction is only able to account for $\\sim$30\\% of the total IR luminosity among the entire sample. ",
        "introduction": "The mid-infrared ($5-40~\\micron$) SED is a complex interplay of broad emission features thought to arise from polycyclic aromatic hydrocarbon (PAH) molecules, silicate absorption features at 9.7 and 18~$\\micron$, and a mid-infrared continuum from very small grains \\citep{lp84,atb85}. Mid-infrared luminosities measured near $\\sim$8~$\\micron$ have been found to correlate with the total infrared (IR; $8-1000$) luminosities of galaxies in the local Universe \\citep{ce01, de02}, which itself is a measure of a galaxy's SFR. While a correlation is found, there does exist a large amount of scatter \\citep{dd05,jd07,la07} and systematic departures for low metallicity systems \\citep{ce05,sm06}. It is therefore uncertain if these empirical relations between mid-infrared and total IR luminosities ($L_{\\rm IR}$) for galaxies in the local Universe are applicable for galaxies at higher redshifts.  Recently, data from deep far-infrared surveys such as the Far-infrared Deep Legacy Survey (FIDEL; PI: M. Dickinson) are confirming the same redshift evolution of the SFR density seen in the mid-infrared and do not show any evolution in the SED of infrared luminous galaxies \\citep[e.g.][]{bm09}. In contrast, \\citet{jr08} have provided evidence that the high redshift lensed mid-infrared selected galaxies might show a factor of $\\sim$2 stronger rest-frame 8~$\\micron$ emission compared to their total IR luminosities.  Using mid-infrared spectroscopy for 22 sources selected at 24~$\\micron$ in the GOODS-N field, along with existing 850~$\\micron$ and additional 70~$\\micron$ imagery obtained as part of FIDEL, we aim to improve our understanding of the star formation and AGN activity within a diverse group of $0.6\\la z \\la2.6$ galaxies. This is done through a more proper estimate of IR luminosities and the ability to decompose these measurements into star-forming and AGN components. For the full study, please see \\citet{ejm09}. ",
        "conclusions": "In the present study we have used observations from the mid-infrared to the submillimeter to properly characterize the IR luminosities for a diverse sample of 22 galaxies spanning a redshift range of $0.6 \\la z \\la 2.6$. In addition, we have used the mid-infrared spectra of these sources to estimate the fractions of their IR luminosities which arise from an AGN. Our conclusions can be summarized as follows:  \\begin{enumerate}  \\item IR ($8-1000~\\micron$) luminosities derived by SED fitting observed 24~$\\micron$ flux densities alone are well matched to those when additional mid-infrared spectroscopy and 16, 70, and 850~$\\micron$ photometry are included in the fits for galaxies having $z \\la 1.4$ and $L_{\\rm IR}^{\\rm 24}$ values typically  $\\la$$3\\times10^{12}~L_{\\sun}$. In contrast, for galaxies lying in a redshift range between 1.4 and 2.6 with $L_{\\rm IR}^{\\rm 24}$ values typically $\\ga$$3\\times10^{12}~L_{\\sun}$, IR luminosities derived by SED template fitting using observed 24~$\\micron$ flux densities alone overestimate the true IR luminosity by a factor of $\\sim$5, on average, compared to fitting all available data. A comparison between the observed mid-infrared spectra with that of the SEDs chosen from fitting 24~$\\micron$ photometry alone and from fitting all available photometric data demonstrates that local high luminosity SED templates show weaker PAH emission by an average factor of $\\sim$5 in this redshift range and do not properly characterize the contribution from PAH emission.   \\item After decomposing the IR luminosity into star forming and AGN components, we find the AGN luminosity to be increasing with increasing difference between the 24~$\\micron$-derived and our best-fit IR luminosities. Such a trend suggests that the AGN power increases with mid-infrared luminosity. However, we also find that the median fraction of the AGN to the difference between the 24~$\\micron$-derived and best-fit IR luminosities is only 16\\% suggesting the AGN power is almost negligible compared to the bolometric correction necessary to properly calibrate the 24~$\\micron$-derived IR luminosities.  \\end{enumerate}"
    },
    "1005/1005.0295_arXiv.txt": {
        "abstract": "{ We study a coupled dark energy--dark matter model in which the energy-momentum exchange is proportional to the Hubble expansion rate. The inclusion of its perturbation is required by gauge invariance. We derive the linear perturbation equations for the gauge invariant energy density contrast and velocity of the coupled fluids, and we determine the initial conditions. The latter turn out to be adiabatic for dark energy,  when assuming adiabatic initial conditions for all the standard fluids. We perform a full Monte Carlo Markov Chain likelihood analysis of the model, using WMAP 7-year data.}  \\preprint{IFT-UAM/CSIC-10-28 \\\\ FTUAM-10-07\\\\ULB-TH/10-15} \\begin{document} ",
        "introduction": "The true substance of dark energy and dark matter is unknown although it should account for about 95\\% of the matter--energy content of our universe today~\\cite{Komatsu:2010fb}. While the couplings of dark fluids to photons and normal matter are severely constrained~\\cite{Carroll:1998zi}, nothing prevents dark matter--dark energy interactions~\\cite{Damour:1990tw,Damour:1990eh,Wetterich:1994bg,Amendola:1999er,Zimdahl:2001ar,Farrar:2003uw,Das:2005yj,Zhang:2005jj,delCampo:2006vv,Bean:2007nx,Olivares:2007rt,Valiviita:2008iv,He:2008si,Gavela:2009cy,Jackson:2009mz,Majerotto:2009np,Valiviita:2009nu,Koyama:2009gd,Boehmer:2009tk}. At the level of the background evolution equations, it is customary to parametrize the coupling between the two dark sectors~\\cite{Kodama:1985bj} as: \\begin{eqnarray} \\label{eq:EOMm} \\dot{\\bar{\\rho}}_{dm}+ 3{\\mathcal{H}}\\bar{\\rho}_{dm} &=&a \\overline{Q}_{dm}\\,,\\\\ \\label{eq:EOMe} \\dot{\\bar{\\rho}}_{de}+ 3 {\\mathcal{H}}\\bar{\\rho}_{de}(1+ w)&=&a \\overline{Q}_{de}\\,, \\end{eqnarray} where  $\\bar\\rho_{dm}, \\, \\bar\\rho_{de}$ denote the dark matter and dark energy energy densities, respectively,  and $\\overline{Q}_{dm}=-\\overline{Q}_{de}$ encodes the coupling between those two dark sectors and drives the energy exchange between them. The dot indicates derivative with respect to the conformal time $d\\tau = dt/a$, with $\\mathcal{H}={\\dot a}/a \\equiv a \\overline{H}$ denoting the background expansion rate, while $w \\equiv w_{de} = \\bar p_{de}/\\bar \\rho_{de}$ stands for the background dark energy equation of state and pressureless dark matter is assumed: $w_{dm}=\\bar p_{dm}/\\bar \\rho_{dm}=0$. From now on, barred quantities are to be considered as the background quantities.   The initial conditions for the several components populating the early universe have been explored to a large extent. They were first analyzed for all cosmic fluids but dark energy (see {\\it e.g.} Ref.~\\cite{Ma:1995ey} and references therein), with the result that adiabatic initial conditions were one possibility. It was also noticed  that the  choice of gauge could be a delicate issue: a safe alternative proposed was to use a gauge invariant formalism \\cite{Bardeen:1980kt,Kodama:1985bj, Mukhanov:1990me,Durrer:2001gq}. The initial conditions for the case of dynamical dark energy as an uncoupled quintessence field have been also derived~\\cite{Viana:1997mt,Dave:2002mn, Malquarti:2002iu,Abramo:2001mv,Kawasaki:2001nx,Perrotta:1998vf,Doran:2003xq}, including a gauge invariant treatment~\\cite{Doran:2003xq}: they turned out to be adiabatic if those for the traditional fluids were adiabatic. Furthermore, the formalism in Ref.~\\cite{Doran:2003xq} has been recently applied to the case of a coupled dark energy-dark matter systems which mimic uncoupled models at early times, both at the background and perturbation levels~\\cite{Majerotto:2009np} for the viable parameter space: as expected, adiabatic initial conditions for dark energy naturally resulted then. Here we consider a different class of dark couplings, not negligible at early times. It is also illustrated that the gauge invariant formalism is  particularly illuminating for the determination of the correct perturbation equations, for a general coupled theory.  The structure of the paper is as follows. In Section~\\ref{sec:gauge-invar-pert}, the notation is set and the gauge invariant equations -at linear order in perturbation theory- for a coupled fluid are derived. In particular, we study in Section~\\ref{sec:coupl-prop-h} the case of a (covariant) dark matter--dark energy interaction proportional to the Hubble rate. In Section~\\ref{sec:initial-conditions}, following the method proposed in Ref.~\\cite{Doran:2003xq}, we derive the corresponding initial conditions for dark energy. Then in Section~\\ref{sec:data-constraints}, we constrain the type of coupled models analyzed, using several data sets. Section~\\ref{sec:conclusion} contains the conclusions. ",
        "conclusions": "\\label{sec:conclusion}  Interacting dark energy-dark matter cosmologies in which the coupling term is proportional to the Hubble  expansion rate are revisited. While in previous works the perturbation in the Hubble expansion rate was neglected, it is illustrated here  how the inclusion of such a term is mandatory to satisfy the gauge invariance of the theory. It also serves as a guide to define a covariant formulation of the dark sector interaction. In this work, the latter has been chosen to be expressed in terms of the expansion rate associated to the total fluid. This choice is however not unique, we could have used the expansion rate of any other fluid. For the case under study, we compute the linear perturbation evolution  using a gauge invariant formalism. After imposing adiabatic initial conditions on the matter and radiation fluids, we find that the initial conditions for the coupled dark energy fluid are also adiabatic. This result is independent of the choice in the covariant formulation of the expansion rate. The new terms arising from the expansion rate perturbation have negligible quantitative impact on the constraints on cosmological parameters previously obtained in the literature.  A new analysis has been performed using the latest WMAP7 data."
    },
    "1005/1005.2290_arXiv.txt": {
        "abstract": "{} {FIR imaging of interacting galaxies allows locating even hidden sites of star formation and measuring of the relative strength of nuclear and extra-nuclear star formation. We want to resolve the star-forming sites in the nearby system of the Antennae.} {Thanks to the unprecedented sharpness and depth of the PACS camera onboard ESA's {{\\em Herschel}} Space Observatory, it is possible for the first time to achieve a complete assessment of individual star-forming knots in the FIR with scan maps at 70, 100, and 160\\,$\\mu$m. We used clump extraction photometry and SED diagnostics to derive the properties related to star-forming activity.} {The PACS 70, 100, and 160\\,$\\mu$m maps trace the knotty structure of the most recent star formation along an arc between the two nuclei in the overlap area. The resolution of the starburst knots and additional multi-wavelength data allow their individual star formation history and state to be analysed. In particular, the brightest knot in the mid-infrared (K1), east of the southern nucleus, exhibits the highest activity by far in terms of dust heating and star formation rate, efficiency, and density. With only 2\\,kpc in diameter, this area has a 10--1000\\,$\\mu$m luminosity, which is as high as that of our Milky Way. It shows the highest deficiency in radio emission in the radio-to-FIR luminosity ratio and a lack of X-ray emission, classifying it as a very young complex. The brightest 100 and 160\\,$\\mu$m emission region (K2), which is close to the collision front and consists of 3 knots, also shows a high star formation density and efficiency and lack of X-ray emission in its most obscured part, but an excess in the radio-to-FIR luminosity ratio. This suggests a young stage, too, but different conditions in its interstellar medium. Our results provide important checkpoints for numerical simulations of interacting galaxies when modelling the star formation and stellar feedback.} {} ",
        "introduction": "We have set up a photometric imaging programme of a sample of nearby interacting galaxies as part of the {\\em Herschel} Guaranteed Time Key Project SHINING\\footnote{http://www.mpe.mpg.de/ir/Research/SHINING/} to trace and resolve the youngest star formation sites that are still embedded in dust, making them strong far-infrared (FIR) emitters. The \\object{Antennae} (\\object{Arp~244}, \\object{NGC~4038}/\\object{NGC~4039}) is a key object in this sample, since it is an interactive system, where strong extranuclear star formation was found previously in the so-called overlap area of the two disks between the two nuclei. Optical images show dark dust lanes interspersed with a few HII regions in this area \\citep[e.g.][]{whitmore95}. It houses several supergiant molecular cloud complexes \\citep{stanford90,wilson00,gao01,schulz07}, which provide a large reservoir for bursts of star formation. Prominent 6 and 20\\,cm radio emission \\citep{hummel86}, which spatially matches the CO emisison peaks, suggests there are ongoing starbursts. Compressed magnetic fields are found towards the northern edge indicating pre-starbursts \\citep{chyzy04}.  Mid-infrared (MIR) observations have discovered embedded starbursts at the southeastern border of the overlap region \\citep{vigroux96,mirabel98,wilson00}, while submm 850 $\\mu$m observations suggest a large amount of cooler dust in the northern overlap region, arguing in favour of relatively less powerful starbursts there \\citep{haas00,zhu03,schulz07}.  In the case of buried starbursts, one expects to see the continuum re-emission of the hiding dust at mid- and far-infrared wavelengths. However, previous FIR space missions have not reached the spatial resolution required to separate the overlap region from the nuclei. {\\em Herschel}-PACS now offers the capability of such spatially resolved observations for the first time. ",
        "conclusions": "New {\\em Herschel}-PACS scan maps of the Antennae at 70, 100, and 160\\,$\\mu$m provide a high spatial resolution complement to earlier MIR maps (ISOCAM, {\\em Spitzer}-MIPS), high-resolution CO (1-0), and radio 6\\,cm maps and thus allow detailed study of the star formation state in individual emission knots.  We confirm that the highest star formation activity takes place in the overlap area concentrated in the two emission complexes K1 and K2. Our star formation diagnostics strongly indicates that these knots differ in evolutionary stage."
    },
    "1005/1005.0589_arXiv.txt": {
        "abstract": "Many global circulation models predict supersonic zonal winds and large vertical shears in the atmospheres of short-period jovian exoplanets.  Using linear analysis and nonlinear local simulations, we investigate hydrodynamic dissipation mechanisms to balance the thermal acceleration of these winds.  The adiabatic Richardson criterion remains a good guide to linear stability, although thermal diffusion allows some modes to violate it at very long wavelengths and very low growth rates.  Nonlinearly, wind speeds saturate at Mach numbers $\\approx 2$ and Richardson numbers $\\lesssim 1/4$ for a broad range of plausible diffusivities and forcing strengths. Turbulence and vertical mixing, though accompanied by weak shocks, dominate the dissipation, which appears to be the outcome of a recurrent Kelvin-Helmholtz instability.  An explicit shear viscosity, as well as thermal diffusivity, is added to ZEUS to capture dissipation outside of shocks.  The wind speed is not monotonic nor single valued for shear viscosities larger than about $10^{-3}$ of the sound speed times the pressure scale height. Coarsening the numerical resolution can also increase the speed. Hence global simulations that are incapable of representing vertical turbulence and shocks, either because of reduced physics or because of limited resolution, may overestimate wind speeds.  We recommend that such simulations include artificial dissipation terms to control the Mach and Richardson numbers and to capture mechanical dissipation as heat. ",
        "introduction": "Strongly irradiated extrasolar giant planets---``hot Jupiters''---are expected to rotate synchronously with their orbits but to have strong longitudinal winds that redistribute heat from the day to the night side.  The efficiency of redistribution is important for direct observables including infrared phase curves, the depth of secondary eclipses in transiting systems (i.e., eclipses of the planet by its star), and planetary spectra.  Turbulence associated with the winds may contribute to chemical mixing of the atmosphere \\citep{Spiegel_Silverio_Burrows09}, and might even inject heat into the convective interior of the planet \\citep{Guillot_Showman02}, thereby perhaps explaining why some of these planets have lower densities than expected by standard evolutionary models at their estimated ages.  In a previous paper \\citep[hereafter Paper I]{Goodman2009}, one of us has argued that these thermally driven winds can be understood as natural heat engines, which convert a fraction of the thermal power into mechanical work: namely, the work expended to accelerate the wind.  As with any other heat engine, the continual production of work must be balanced by mechanical dissipation, else the kinetic energy in the winds would grow without bound.  Paper I offered a brief discussion of possible dissipation mechanisms.  Here we explore those mechanisms in greater detail.  The goal is to lay the foundation for subgrid models of the dissipative process suitable for use by global circulation codes.  The outline of this paper is as follows.  \\S\\ref{sec:overview} gives an overview of the dissipation mechanisms we consider. \\S\\ref{sec:linear} presents a linear analysis of hydrodynamic Kelvin-Helmholtz instabilities of thermally stratified shear flows.  We show that such flows can be destabilized by thermal diffusivity even at Richardson numbers greater than the well-known critical value $Ri_{\\rm crit}=1/4$.  Under conditions relevant to exoplanet winds, however, we estimate that the associated growth rates are too slow and the longitudinal wavelengths too long to be important for dissipation, unless $Ri<1/4$, which requires transsonic flow unless the vertical width of the shear layer is small compared to a pressure scale height.  Hence shocks may be as important as shear-driven turbulence for dissipation.  To investigate this, as described in \\S\\ref{sec:nonlinear}, we have used the ZEUS code in two dimensions to simulate a part of the atmosphere with horizontal and vertical dimensions comparable to the pressure scale height.  Thermal diffusion and viscosity have been added to the code, and thermal driving of the wind is simulated by a horizontal body force with nonzero curl.  We study the velocity and dissipation rate of the wind in a statistical steady state as a function of the strength of the driving compared to the acceleration of gravity. \\S\\ref{sec:discussion} discusses our results in the context of previous work on winds in jovian planets both within and beyond the solar system.  \\S\\ref{sec:conclusions} summarizes our main conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We have investigated the balance between horizontal forcing and mechanical dissipation in model shear flows designed to imitate thermally driven winds of jovian exoplanets. Our main conclusions are as follows. \\begin{enumerate} \\item Following Paper I, some form of mechanical dissipation is required to offset the production of mechanical energy by the longitudinal entropy gradient.   Plausible dissipation mechanisms include shocks, turbulence, and MHD drag.  \\item Entropy stratification has a stabilizing influence on Kelvin-Helmholtz (KH) modes, but  thermal diffusivity ($\\chi$) tends to undercut this influence to a degree that can be quantified by the Peclet number $Pe\\equiv\\cs H_p/\\chi$ based on the sound speed ($\\cs$) and pressure scale height ($H_p$).  We estimate $Pe\\sim O(10^3)$ at the altitude where the thermal timescale is comparable to the circumferential flow time at Mach~1, but  $Pe\\lesssim O(1)$ near the infrared photosphere.  \\item Linear analysis indicates that for $Pe\\gtrsim 10$, thermal diffusion has little effect on the stability of KH modes whose horizontal wavelength is comparable to the vertical thickness of the shear layer; the usual adiabatic Richardson criterion $Ri\\equiv N^2/(\\partial U/\\partial z)^2>1/4$ is sufficient for the stability of these short-wavelength modes.  For typical conditions on the day sides of hot Jupiters, $Ri<1/4$ implies transsonic or supersonic winds, presuming that the shear layer is no narrower than a pressure scale height.  \\item Nevertheless, at any finite $Pe$ and $Ri$, inviscid flows with an inflection point in the shear profile are unstable at sufficiently long horizontal wavelengths, at least when the flow is confined within a channel of finite vertical extent.  We doubt that these long-wavelength KH modes are important for hot-Jupiter winds because their growth rates are very small (inversely proportional to wavelength) and because the planetary circumference is not large enough compared to $H_p$ to accommodate them at $Pe\\gtrsim 10^3$.  \\item We estimate a turbulent diffusivity for composition in hot Jupiter atmospheres.  For $Pe\\approx 100$, we find $K_{\\rm t}/c_sH_p\\approx 10^{-3}$, depending on the strength of the horizontal forcing.  However, these values pertain to regions of maximum vertical shear; turbulent diffusivities may be much smaller elsewhere.  \\item Nonlinear, two-dimensional (azimuth and depth), horizontally forced local simulations saturate at time-averaged Mach numbers $\\sim 2$  when the Reynolds number $Re\\equiv \\cs H_p/\\nu\\gtrsim 10^3$.  Given the widths of our shear layers ($\\sim 2H_p$) and background stratification, this corresponds to minimum Richardson numbers slightly below the adiabatic critical value of $1/4$. In this regime, which is relevant to hot Jupiters, we find that the saturated wind speed depends only weakly on $Pe$ and on the strength of the forcing.  \\item Dissipation in the high-$Re$ regime is due mainly to the work expended to overturn the stratification rather than shocks, though weak shocks are present.  It is highly intermittent, at least in our two-dimensional calculations, and appears to be triggered by a recurrent linear instability of Kelvin-Helmholtz type.  \\item At $Re\\lesssim 10^3$ or at low numerical resolution (significantly fewer than 10 cells per pressure scale height), multiple sequences of solutions appear in overlapping ranges of Reynolds number, differing in their wind speed and predominant dissipation mechanisms.  The fastest have Mach numbers at least twice as large as what we believe to be the correct inviscid values for our forcing strengths.  \\end{enumerate}  The last three points suggest that some global simulations of hot Jupiters may have overestimated wind speeds due to incomplete physics (e.g. neglect of vertical accelerations and sound waves) and/or insufficient spatial resolution.  Pending sufficient computational resources to resolve the physical dissipation mechanisms directly in global simulations, we suggest that artificial dissipative terms be added to the equations of motion so as to prevent strongly supersonic shears and small Richardson numbers.  The appropriate form for these terms will depend upon the equation sets and algorithms used, but presumably it is desirable that they conserve energy and momentum, and probably that they have a sharp threshold in Mach number or Richardson number."
    },
    "1005/1005.3858_arXiv.txt": {
        "abstract": "The fragile structure of chondritic-porous interplanetary dust particles (CP- IDPs) and their minimal parent-body alteration have led researchers to believe these particles originate in comets rather than asteroids where aqueous and thermal alteration have occurred. The solar elemental abundances and atmospheric entry speed of CP-IDPs also suggest a cometary origin. With the return of the Stardust samples from Jupiter-family comet 81P/Wild 2, this hypothesis can be tested. We have measured the Fe oxidation state of 15 CP-IDPs and 194 Stardust fragments using a synchrotron-based x-ray microprobe. We analyzed $\\sim$300 nanograms of Wild 2 material -- three orders of magnitude more material than other analyses comparing Wild 2 and CP-IDPs. The Fe oxidation state of these two samples of material are $>$2$\\sigma$ different: the CP-IDPs are more oxidized than the Wild 2 grains. We conclude that comet Wild 2 contains material that formed at a lower oxygen fugacity than the parent body, or parent bodies, of CP-IDPs. If all Jupiter-family comets are similar, they do not appear to be consistent with the origin of CP-IDPs. However, comets that formed from a different mix of nebular material and are more oxidized than Wild 2 could be the source of CP-IDPs. ",
        "introduction": "\\label{intro} Chondritic-porous interplanetary dust particles (CP-IDPs), collected in the stratosphere by high-altitude aircraft, are widely thought to originate in comets \\citep{Schramm:1989p3433,Bradley:1994p2672}. The collisional history of asteroids makes it unlikely that these particles (fragile, porous aggregates of small grains) are asteroidal \\citep{Brownlee:1985p2339}.  The atmospheric entry speed of CP-IDPs, as deduced from thermal release profiles of solar wind He \\citep{Nier:1992p3404}, is consistent with cometary, rather than asteroidal, orbits \\citep{Brownlee:1995p2673}. Infrared and electron beam studies of IDPs show that those consisting of mostly anhydrous minerals are usually chondritic-porous, whereas IDPs rich in phyllosilicates are usually chondritic-smooth \\citep{Sandford:1985p3456, Bradley:1986p3458}. The presence of GEMS (glass with embedded metal and sulfides) in IDPs indicates a lack of thermal alteration \\citep{Bradley:1994p2672}. CP-IDPs, therefore, show very little parent-body alteration and must originate from either anhydrous objects or hydrous objects that have been kept at very low temperature -- again, consistent with cometary origin. Nevertheless, a cometary origin of CP-IDPs is not universally accepted \\citep{Flynn:1992p3174, Thomas:1995p3175}. With the return of samples from comet 81P/Wild 2 by NASA's Stardust mission \\citep{Brownlee:2006p1376}, it is now possible to compare CP-IDPs with material from this Jupiter-family comet. We seek to prove or disprove, with a known level of confidence, the hypothesis that CP-IDPs originate from parent bodies with the composition of comet Wild 2.  The oxidation state of Fe is a clear mineralogic indicator of the oxidation state of meteorites \\citep{Rubin:1988p2665}. Meteorite groups are in fact distinguished from each other by their differing Fe oxidation states, as Urey and Craig first reported more than fifty years ago \\citep{Urey:1953p1396}. Parent body processing can drastically change the Fe oxidation state of some of the meteorite groups, obscuring information about the original oxidation state the material acquired in the solar nebula \\citep{Rubin:1988p2665}. Carbonates and Fe-bearing crystalline silicates, possible products of aqueous alteration on asteroids \\citep{Krot:1995p3169}, are found alongside anhydrous minerals in the same Stardust track \\citep{Flynn:2008p2675}, and therefore co-existed in one large aggregate in comet Wild 2 . No phyllosilicates have been unambiguously identified in the Stardust samples \\citep{Zolensky:2008p1420}.  With these and other pieces of evidence, \\citet{Wooden:2008p2674} argues that instead of selective aqueous alteration on submicron scales in the comet itself, grains which formed in different regions of the solar nebula under varying reduction-oxidation conditions (e.g. Mg- and Fe-rich crystalline silicates) migrated, aggregated, and formed comet Wild 2. Likewise, the oxidation state of anhydrous CP-IDPs was unaffected by parent body processing \\citep{Schramm:1989p3433,Zolensky:1995p1397}, so a comparison of the Fe oxidation state of comet Wild 2 with anhydrous CP-IDPs yields insight into the nebular environment in which they formed.  We report on the oxidation state of Fe in 15 anhydrous CP-IDPs, and directly compare these measurements to the Fe oxidation state of comet Wild 2, deduced from 194 Stardust fragments in 11 aerogel tracks. ",
        "conclusions": "The presence of Fe in its different oxidation states is one of the clearest mineralogic indicators of the oxidation state of meteoritic material \\citep{Rubin:1988p2665}. As discussed earlier, the Wild 2 samples and CP-IDPs are unlikely to have experienced much parent-body processing that would drastically change their oxidation state, and therefore a measurement of the Fe oxidation state of these two materials is a probe into the oxygen fugacity of the environment in which they formed. Figure \\ref{figure:ternary} shows that comet Wild 2 and CP-IDPs differ in their Fe oxidation state at more than the 2$\\sigma$ level.  \\citet{Ishii:2008p962} reported several differences between CP-IDPs and the Wild 2 material: the Stardust material lacks GEMS (glass with embedded metal and sulfides) and also lacks elongated enstatite whiskers and platelets, all of which are present in CP-IDPs. They conclude from these transmission electron microscopy observations that the Wild 2 material is dissimilar to CP-IDPs and more closely resembles asteroidal material. Our results, analyzing $\\sim$1000$\\times$ more material, indicate that comet Wild 2 differs at more than the 2$\\sigma$ confidence level from CP-IDPs, and Wild 2 is also distinct from meteorites. CP-IDPs are 2$\\sigma$ inconsistent with all meteorite groups except CI and R.  The significant difference between the Fe oxidation state of Wild 2 and CP-IDPs is that a significant fraction of Fe in Wild 2 exists as Fe$^{0}$ whereas CP-IDPs contain little Fe$^{0}$ compared to Fe sulfide and oxidized Fe. We can compare these results only qualitatively to, for example, TEM observations of mineral phases present in CP-IDPs: the Fe XANES technique used in this work can determine how many Fe atoms in the entire CP-IDP are in a given Fe oxidation state, whereas the the entire volume of the IDP is not always examined by TEM, and the amount of material in a given phase is seldom recorded. \\citet{MacKinnon:1987p2660} tabulate a collection of phases observed in 30 CP-IDPs. Of the 30, Fe-bearing mineral phases are reported in 23 of the CP-IDPs: 21 contain oxidized Fe, 15 contain Fe sulfide, and 4 contain Fe$^{0}$. Qualitatively, this is consistent with the relative paucity of Fe$^{0}$ we see by the Fe XANES technique employed in this work.  \\subsection{Selection bias in CP-IDP collection?}  Since Fe metal is denser than Fe sulfides or oxidized Fe, particles containing substantial amounts of Fe metal may have a shorter residence time in the stratosphere and are therefore less likely to be collected. Particles containing Fe metal are also more likely to vaporize than less-dense particles \\citep{Flynn:2008p3176}.  Can the effect of these biases skew the composition of collected IDPs such that our observations of the 15 IDPs studied in this work are actually consistent with IDPs of Wild 2 composition before atmospheric entry? To answer this question, we recalculate the Fe oxidation for the 11 Stardust tracks from comet Wild 2 as if they were IDPs, accounting for the effects of atmospheric heating and residence time in the stratosphere.  First, we first calculate the residence time of 11 IDPs in the range of 15--25 km of altitude from the formulation given in \\citet{Kasten:1968p3375} and \\citet{Rietmeijer:1993p3303}. The IDPs have the Wild 2 Fe composition as given by the Stardust tracks in Table \\ref{table:fefrac}. The speed of the particle falling through the atmosphere is proportional to its density, so the less-dense particles are more likely to be collected than the more-dense, Fe-metal-rich particles. Secondly, we consider the atmospheric entry heating of IDPs as described by \\citet{Fraundorf:1980p3177}. We calculate the probability that an IDP is destroyed by heating to a temperature above the melting point of Fe metal, 1811$^{\\circ}$ K.  The probability of collecting IDP $i$ of all $j$ types of IDPs can be expressed simply as:  \\begin{equation} P_i=\\frac{\\left(N_i-N_iV_i\\right)T_i}{\\displaystyle\\sum_j{\\left(N_j-N_jV_j\\right)T_j}} \\label{entryeqn} \\end{equation}  where $N_i$ is the number of IDPs of type $i$, $V_i$ is the fraction of IDPs $i$ that vaporize, and $T_i$ is the residence time of IDP $i$ in the stratospheric-collection region.  For each of the 11 IDPs of Wild 2 composition, we calculate the probability of collection. Then we calculate the mean and 2$\\sigma$ confidence limits of the Fe oxidation state of these 11 particles using the method described in Section \\ref{confidencelimits}, but here each IDP is weighted by its probability of collection relative to the others.  The amount of Fe metal in IDPs of Wild 2 composition, taking into account atmospheric heating and residence time in the stratosphere, is $<$1\\% less than the abundance of Fe metal in the IDPs before atmospheric entry. This cannot explain the observed lack of Fe metal in CP-IDPs compared to comet Wild 2, so we conclude that these possible selection biases in IDP collection are insignificant.  \\subsection{Grain sizes in CP-IDPs and Stardust Fragments} CP-IDPs are aggregates of mostly submicron grains whereas the Stardust fragments contain many mono-mineralic grains larger than a micron. Thermal modification or destruction of submicron grains in the bulb \\citep{Leroux:2009p3050,Brownlee:2006p1376} precludes the possibility of comparing the Fe oxidation state of CP-IDPs and Stardust at the same grain size. It may be the case that cometary metal exists only in larger grains and is not present in the submicron components that make up CP-IDPs. While we cannot rule out this possibility with the measurements presented here, we can compare comet Wild 2 and CP-IDP material by looking at smaller grain sizes of Wild 2 (What is the Fe$^{0}$ fraction of the smaller Stardust fragments we analyzed?) and larger grain sizes of CP-IDPs (Do larger mineral grains of CP-IDPs exist and what is their Fe oxidation state?).  For each Stardust fragment we analyzed, we know its approximate, relative Fe mass by the absorption edge jump across the Fe K-edge. We compute Fe$^{0}/\\sum\\,$Fe for all Stardust fragments less than a given Fe mass, shown in Figure \\ref{figure:FeMassFrac}. The Fe$^{0}$ fraction increases for the intermediate-mass particles in this study, but is still inconsistent with the CP-IDPs over the entire Fe-mass range of Stardust fragments we analyzed (using the particle-based bootstrap uncertainty calculation).  \\begin{figure} \\centering \\includegraphics[width=\\textwidth]{fefracsize} \\caption{Fe$^{0}/\\sum\\,$Fe for all Stardust fragments less than a given Fe mass (computed by the Fe K-edge XANES jump). The 2$\\sigma$ confidence interval for each set of Stardust fragments less than a given mass is shown by the red shaded region. The green region indicates the 2$\\sigma$ confidence interval for the CP-IDPs measured in this work (CP-IDPs consist of mostly submicron grains--smaller grain size than the Stardust fragments analyzed in this work). Both the Stardust and CP-IDP uncertainties shown here are calculated with the particle-based approach described in Section \\ref{confidencelimits}.} \\label{figure:FeMassFrac} \\end{figure}  Stratospheric collection yields large ($>$5$\\mu$m) mono-mineralic grains of non-chondritic composition in addition to IDPs. Fine-grained chondritic material (reminiscent of small bits of CP-IDPs) is sometimes found adhering to the surfaces of these grains, indicating they originate from the same parent body or bodies as CP-IDPs \\citep{Flynn:2007p3428,Flynn:2008p3176}. The sizes of these grains are closer to the range of grain sizes of the Stardust fragments we analyzed. The mono-mineralic IDP grains consist predominantly of olivine, pyroxene, and sulfide. They probably do not contain a significant amount Fe metal. Therefore, these larger mono-mineralic IDP grains which likely sample the same parent body as the CP-IDPs, probably do not contain enough Fe$^{0}$ to be consistent with the comet Wild 2 material.  If a cosmic dust particle is highly enriched in Fe metal to the exclusion of most other phases, its chemical composition would not be close to chondritic and therefore the particle would not be considered a CP-IDP. Similarly, a particle made mostly of Fe metal may not appear fluffy, and so we may not select it from the Cosmic Dust Catalogs for inclusion in a study such as this. There is an inherent bias in the selection of CP-IDPs for comparison with other samples of matter. Stratospheric-collected particles with high Fe metal abundances have been classified as ``Cosmic\", but the fraction of the collected Fe metal particles which are of true extraterrestrial origin or possibly cometary has not been established. However, if Fe metal were present as submicron grains at the 10\\% level, a particle could still have approximately chondritic composition and appear fluffy, like a CP-IDP. This, we have shown, is not the case. For this type of selection bias to significantly affect our results, the Fe metal distribution in IDPs would have to be very heterogenous: the Fe metal would need to be sequestered into particles which contain enough Fe metal to appear either non-chondritic or non-porous.  Our analysis compares sub-micron-grained CP-IDPs with larger-grained Stardust fragments, but as the Stardust fragments show more Fe$^{0}$ than CP-IDPs over the entire range of particle sizes we measured, and large mono-mineralic grains which likely originated from the same source material as IDPs also show a dearth of Fe$^{0}$, we conclude it is likely that the CP-IDPs contain less Fe$^{0}$ than the Wild 2 samples at comparable grain sizes. We acknowledge a possible selection bias in choosing chondritic and porous IDPs for this study, though our results are valid under the assumption that Fe metal is not distributed heterogeneously, with scarce, very Fe-metal-rich particles containing the Fe metal missing in CP-IDPs.  \\subsection{Fe metal in Stardust samples} The Fe metal we observe in the Wild 2 samples is unlikely to be a capture effect, as discussed earlier. \\citet{Zolensky:2006p1378} report that the high abundance of Ni in metal droplets in Stardust tracks shows that there must have been metal intrinsic to the comet. Much of the Fe metal we measure exists in large particles in the terminus of the track. One of the terminal particles in Track 77 was observed in TEM to be an Fe metal grain encased in olivine (D. Joswiak, private communication), reminiscent of Fe metal encapsulated in silicate grains in unequilibrated chondrites \\citep{Rambaldi:1980p2671}. The terminal particle of Track 41 is a large piece of Fe metal (as shown by the Fe XANES, and an x-ray diffraction match to kamacite and taenite) with a bit of Fe sulfide and Fe$^{2+}$ (as seen in the chemical map of the particle, Figure \\ref{figure:simeio}). The Fe metal in large terminal particles like these almost certainly existed as metal in the comet, which makes Wild 2 distinct from CP-IDPs.  \\begin{figure} \\centering \\includegraphics[width=\\textwidth]{simeio_chemical_map_figure_blue_gamma} \\caption{Iron chemical map \\citep{Ogliore:2008p3135} of the the large terminal particle in Track 41, ``Simeio\", in two orientations that differ by a 90$^{\\circ}$ rotation. Simeio is mostly a large mass of Fe$^{0}$. The mineral phase of this material was identified as kamacite and taenite by x-ray diffraction (XRD). Iron sulfide is present in small amounts throughout the particle, as well as both Fe$^{2+}$ and sulfide on the particle perimeter in both orientations, possibly indicating a thin, partial shell of sulfide and Fe$^{2+}$. } \\label{figure:simeio} \\end{figure}  \\subsection{Nebular mixing, CP-IDPs, and comet Wild 2} The oxidation state of Fe can be used as an indicator of the oxygen fugacity in which the grain formed \\citep{Rubin:1988p2665}. These metal grains likely formed in the high-temperature, low-O-fugacity conditions of the inner region of the early solar nebula (a reducing environment \\citep{Wooden:2008p2674}). Refractory inclusions from the inner nebula have been identified in the Stardust samples \\citep{Simon:2008p2065}. Our Fe XANES survey of Stardust tracks imply that a fraction between 0.5 and 0.65 (2$\\sigma$) of Wild 2 material came from the inner nebula \\citep{Westphal:2009p1874}, suggesting that there was large-scale mixing in the solar nebula \\citep{Ogliore:2009p2695}. Stardust olivines show a broad distribution of Mg/Fe composition \\citep{Zolensky:2006p1378} which indicates the comet consists of grains that formed in a wide range of reducing/oxidizing conditions \\citep{Rubin:1988p2665}. Therefore, it is not unexpected that the oxidation state of Fe in Wild 2 should reflect both the inner nebula (Fe metal grains) and the outer nebula (Fe-rich silicates \\citep{Wooden:2008p2674}).  The fact that CP-IDPs show much less Fe metal than Wild 2 indicate that these particles did not originate from a body, or bodies, that sampled the regions of low oxygen fugacity that Wild 2 did. Comets can contain different amounts of low and high oxygen fugacity material, as \\citet{Wooden:2008p2674} states: ``Hence, the efficiency of incorporation of low vs. high oxygen fugacity products into cometary nuclei depends on the interplay between the duration of low vs. high water vapor content, the time-dependent outward mass transport rate, and the distance(s), range of radial migration, and duration of cometesimal accretion and nuclei accumulation\". CP-IDPs could be derived from comets that did not sample the highly mixed nebula as Wild 2 did, perhaps one that formed in a different time or place.  Our results are inconsistent with the hypothesis that CP-IDPs originate from parent bodies with the composition of comet Wild 2 at $>$2$\\sigma$ confidence level. We have not, however, disproved the hypothesis that CP-IDPs originate from Jupiter-family comets because we do not know if all Jupiter-family comets have an Fe oxidation state that is consistent with comet Wild 2.  Comet Tempel 1, the target of NASA's Deep Impact mission, shows significant differences from Wild 2 in composition \\citep{Brownlee:2006p1376} and morphology of the cometary nucleus \\citep{Thomas:2007p3161}. These differences could be due to the longer period of time Tempel 1 spent in the inner solar system, the different regions from which the samples originated, the different size-range of analyzed material, and many other complicating factors. However, it is also possible that the two comets are composed of fundamentally different material that formed in different places and at different times in the solar nebula. From orbital dynamics, \\citet{Liou:1996p3165} show that dust from Tempel 1 could contribute significantly to the population of low-entry-speed CP-IDPs collected at Earth. This work has shown that a cometary source of CP-IDPs must contain less inner nebula material (where the oxygen fugacity is low) than comet Wild 2--we do not rule out a CP-IDP source from comet Tempel 1 or the numerous other comets of unknown composition."
    },
    "1005/1005.3543_arXiv.txt": {
        "abstract": "{We present five band imaging of the Vega debris disc obtained using the \\emph{Herschel Space Observatory}.  These data span a wavelength range of 70--500\\,$\\mu$m with full-width half-maximum angular resolutions of 5.6-36.9\\arcsec.  The disc is well resolved in all bands, with the ring structure visible at 70 and 160\\,$\\mu$m.  Radial profiles of the disc surface brightness are produced, and a disc radius of 11\\arcsec\\ ($\\sim 85$\\,AU) is determined.  The disc is seen to have a smooth structure thoughout the entire wavelength range, suggesting that the disc is in a steady state, rather than being an ephemeral structure caused by the recent collision of two large planetesimals.} ",
        "introduction": "Debris discs, of which the \\object{Vega} ($\\alpha$ Lyrae) disc is the archetype, are characterised as discs of dusty material generated by the collision of planetesimals in belts surrounding main sequence stars.  The ages of the stars which exhibit these discs ($\\sim$ 350\\,Myr in the case of Vega; \\citealt{Song2000}) precludes the possibility for this dust to be primordial, as the time scale to remove such dust is $\\la$10\\,Myr \\citep{Backman1993,Wyatt2008}.  The debris disc around Vega was first detected by \\cite{Aumann1984} as an infrared excess using the Infrared Astronomical Satellite (\\emph{IRAS}; \\citealt{IRAS1984}), and has been extensively studied in the infrared and submillimetre over the subsequent 25 years \\citep[e.g.][]{Holland1998, Wilner2002, Su2005, Marsh2006}.  The appearance of the disc has been found to vary significantly across this wavelength regime, changing from a smooth axisymmetric structure in the infrared \\citep[hereafter S05]{Su2005}, to a structure in the submillimetre, wherein the majority of the emission lies in two discrete clumps \\citep{Holland1998}.  In order to understand the reason for the variation in structure with wavelength it is important to first understand the origin of the clumps seen in the submillimetre.  The recent collision of two massive planetesimals is one option, however, given the age of Vega, the statistical likelihood of this occurring with two bodies of sufficient mass to explain the submillimetre observations is low \\citep{Wyatt2002}.  A more favourable alternative, first proposed by \\cite{Wilner2002} and modelled by \\citet[hereafter W06]{Wyatt2006} and \\cite{Reche2008}, is that the clumps are dust grains trapped in resonance with a planet near to the disc.  In this scenario the large dust grains (larger than a few mm) are trapped in these resonances, while smaller intermediate sized grains (a few $\\mu$m--mm), having been perturbed by radiation pressure, have a more uniform distribution in the disc.  Recent analysis and modelling of \\emph{Spitzer} \\citep{Werner2004} mid-infrared data have reached contradictory conclusions. S05 find the disc to be ephemeral; in this scenario  the disc is the result of a recent massive collision of planetesimals, and the subsequent collisional cascade.  This results in a high mass of very small grains (less than a few $\\mu$m) which are blown out of the system by radiation pressure immediately upon creation, resulting in the large disc extent observed.  Conversely, \\cite{Muller2010} succeed in reproducing the surface brightness radial profile using intermediate size grains in elliptical orbits around the parent planetesimal ring, and therefore conclude that it is consistent with a steady-state model.  In the steady-state model, dust that is destroyed, either by being drawn in to the star due to Poynting-Robertson drag or blown out of the disc by radiation pressure, is continuously replenished by a steady collisional cascade within the planetesimal belt.  If the small blown-out grains are the origin of the emission observed in the mid-infrared then W06 predicts that spiral features, emanating from the submillimetre clumps, should be visible with high-resolution imaging; a smooth structure would support the steady-state model.  In this paper we present five-band far-infrared imaging of the Vega debris disc obtained with the \\emph{Herschel} \\citep{Pilbratt2010} Photodetector Array Camera and Spectrometer \\citep[PACS;][]{Poglitsch2010} and Spectral and Photometric Imaging Receiver \\citep[SPIRE;][]{Griffin2010}.  We discuss the initial analysis and disc parameterisation, and relate these results to the ephemeral and steady-state disc models. In Section 2 we present the \\emph{Herschel} data, outline the processing performed, and analyse the disc structure and properties.  These data are then compared with results from the recent \\emph{Spitzer} observations (S05) and disc modelling of W06 and \\cite{Muller2010}, with our conclusions summarised in Section 3. ",
        "conclusions": "The structure observed in the \\emph{Herschel} data shows no sign of clumps. There are also no visible spiral arm features, predicted by W06, if the disc emission at mid-to far-infrared wavelengths is dominated by small blown-out dust grains (W06 Fig.~3; right hand panels, $\\beta=$1--10).  The smooth structure observed is most consistent with the steady-state model, wherein the emission is dominated by intermediate size dust grains in elliptical orbits about the parent planetesimal belt.  This model was found to simultaneously give good agreement to the data in all bands, which is unexpected, as the more distant grains should have a lower temperature, and suggests that the mean grain size decreases with distance from the star.  This is in-keeping with observational data which shows a larger disc at shorter wavelengths.  Full modelling of the radial grain size distribution will be presented in Sibthorpe (2010, in prep). \\begin{enumerate} \\item We presented resolved images of the Vega debris disc system in five bands ranging from 70--500\\,$\\mu$m obtained using the \\emph{Herschel} PACS and SPIRE instruments. \\item The peak surface brightness of the dust disc was identified at 70 and 160\\,$\\mu$m at a radius of 11\\arcsec\\ (85 AU). \\item The surface brightness profile was found to be well fit in the outer disc by a $\\log_{10}(S_{\\nu}) \\propto -0.63r$ distribution, with a different scale factor at each band.  The inner profile ($r \\le r_{0}$) was likewise modelled, with a Gaussian profile of FWHM = 20\\arcsec\\ found to provide a good fit.  The change in surface brightness distribution, occuring at a radius of $\\sim$14\\arcsec\\ ($\\sim$109 AU), is used to observationally define the distinction between the inner and outer disc. This model was found to simultaneously give good agreement to the data in all bands. \\item The structure of the disc was found to be smooth, with no clumpy structure to the sensitivity limit of these data. \\item While these data cannot preclude the option that the Vega disc is the result of a large planetesimal collision, making it ephemeral in nature, these data support the hypothesis that the Vega disc is steady-state in nature. \\end{enumerate}"
    },
    "1005/1005.1770_arXiv.txt": {
        "abstract": "One of the main tasks for present and future dark energy surveys is to determine whether the dark energy is dynamical or not. To illustrate this from data, it is commonly used to parameterize the dark energy equation of state $\\omega$ as several piecewise constant $\\omega_{i}$s using the principal component analysis (PCA) method over finite redshift bins. We show that there is only $j-1$ free parameters $\\omega_{i}$s if we choose the redshift as $j$ bins. Without this constrain, one obtains the inconsistent results from the data analysis. Furthermore, if $\\omega$ decreases with non-negligible ratio as $z$ does, then PCA fails to reproduce the original behavior of $\\omega$. Also, time varying $\\omega$ can be confused with the incorrect value of constant one when the decreasing (or increasing) ratio of $\\omega$ is small but not negligible. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.3296_arXiv.txt": {
        "abstract": "We report on the results from \\halpha\\ imaging observations of the eastern limb of Tycho's supernova remnant (SN1572) using the Wide Field Planetary Camera-2 on the Hubble Space Telescope. We resolve the detailed structure of the fast, collisionless shock wave into a delicate structure of nearly edge-on filaments. We find a gradual increase of \\halpha\\ intensity just ahead of the shock front, which we interpret as emission from the thin ($\\sim1\\arcsec$) shock precursor. We find that a significant amount of the \\halpha\\ emission comes from the precursor and that this could affect the amount of temperature equilibration derived from the observed flux ratio of the broad and narrow \\halpha\\ components. The observed \\halpha\\ emission profiles are fit using simple precursor models, and we discuss the relevant parameters. We suggest that the precursor is likely due to cosmic rays and discuss the efficiency of cosmic ray acceleration at this position. ",
        "introduction": "\\label{sec:intro}     The shock transition in fast astrophysical shocks is intrinsically a ``collisionless'' process, and energy is dissipated via plasma turbulence and/or electromagnetic fields.  An important consequence of the collisionless nature of the shocks is cosmic ray acceleration \\citep[e.g.,][]{1987PhR...154....1B}. While there is increasing evidence of cosmic ray acceleration in supernova remnants, the details of the process are still not well understood, and the question of whether supernova remnants (SNRs) are the primary acceleration sites of Galactic cosmic rays is still open \\citep{2009Natur.460..701B}.     Cosmic ray acceleration models require a precursor in which accelerated particles can scatter back to the postshock region for further acceleration. Observations of the cosmic ray precursor can constrain the two key parameters of acceleration models; the diffusion coefficient and the injection efficiency \\citep{1987PhR...154....1B,1988ApJ...333..198B}. The Balmer-dominated filaments that are produced when fast SNR shocks propagate into partially neutral gas are potential sites where such a cosmic ray precursors can be observed. Most of the Balmer emission comes from a very narrow zone behind the shock, where the hydrogen atoms swept up by the shock are excited before they are ionized \\citep{1978ApJ...225L..27C,1980ApJ...235..186C}.  However, in the presence of the precursor, additional Balmer emission is expected from the precursor region where the preshock gas is compressed and heated. Using long-slit \\halpha\\ spectroscopy along the shock normal of a Balmer filament in Tycho's SNR, \\citet[][Lee07 hereafter]{2007ApJ...659L.133L} found that there is an increase of the \\halpha\\ narrow component intensity in a small region ($\\sim 0.4\\arcsec$) ahead of the shock front, which they proposed as potential emission from the precursor. However, the angular resolution of the observation by Lee07 is $\\sim 0.5\\arcsec$ and their results needed to be verified with high resolution observations.    In this \\emph{Letter}, we report \\halpha\\ imaging observations of Balmer-dominated filaments in Tycho's SNR using the Wide Field and Planetary Camera 2 (WFPC2) on the Hubble Space Telescope (HST), which resolves the detailed structure of the shock. \\S~\\ref{sec:observations} presents the observations and reports the detection of the precursor. The precursor is modeled in \\S~\\ref{sec:model} and its characteristics are discussed in \\S~\\ref{sec:analysis}. Finally in \\S~\\ref{sec:summary}, we discuss the efficiency of cosmic ray acceleration in this region.    \\begin{figure}[b] \\plotone{f1.eps} \\label{fig:kpno-hst} \\caption{ Hubble Space Telescope image of the Balmer-dominated filaments in the northeastern part of Tycho's SNR.  The WFPC2 detector with F656N filter was used. The inset shows an \\halpha\\ image toward Tycho's SNR from a ground-based telescope (KPNO). The area observed by HST is marked. } \\end{figure} ",
        "conclusions": "The existence of the cosmic ray precursor does not necessarily imply efficient cosmic ray acceleration. For the Balmer-dominated filaments to be observable, some neutral hydrogen needs to survive ionization in the precursor. As efficient cosmic ray acceleration tends to make the precursor wider and hotter, the Balmer-dominated filaments may not trace the shocks having efficient cosmic ray acceleration \\citep[cf.][]{2009Sci...325..719H}. For the shock studied in this paper, a more direct suggestion of inefficient acceleration can be inferred from the results of Lee07. From the difference in radial velocity of the preshock gas and the \\halpha\\ narrow component emitted in the postshock region, they estimated the amount of gas deceleration in the precursor. The preshock gas is decelerated in the precursor due to the gradient of the cosmic ray pressure by about a few hundred \\kms\\ based on the radial velocity measurements. The value is not sensitive to the assumed distance, and is significantly smaller than the measured line width of the broad component.  % In this shock, the thermal pressure of the ordinary gas still dominates over the cosmic ray pressure and the acceleration is not likely to be very efficient.    Throughout the discussion, we have assumed that the distance to Tycho's SNR is 2.1 kpc. The measurement is based on the estimated proton temperature (from the observed line width of the broad component) and the optical proper motion, assuming no cosmic ray acceleration at the shock. The cosmic ray acceleration, if significant, can effectively reduce the postshock temperature \\citep[see][and references therein]{2009Sci...325..719H}; thus the distance may have been underestimated. However, as has been discussed above, the optical observations are consistent with cosmic ray acceleration not being very efficient in this region, so the distance of 2.1 kpc, derived assuming no cosmic ray acceleration, remains a reasonable value.   Using Chandra observations, \\citet{2005ApJ...634..376W} found the locations of the shock front (SF) and the contact discontinuity (CD) along the boundary of the remnant. They interpreted the small separation between the two as an indication of efficient cosmic ray acceleration.  The region of knot g is one of the regions where the SF-CD separation is smallest (except those regions of ejecta protrusion). A simple extrapolation of their argument will lead to the most efficient cosmic ray acceleration in this region, being inconsistent with our results.  However, the region around knot g is where the remnant could be interacting with dense ambient clouds \\citep{1997ApJ...491..816R,2004ApJ...605L.113L}.  Thus, the small SF-CD separation in this region could be due to a recent encounter of the shock with the dense ambient gas, instead of efficient cosmic ray acceleration.   In conclusion, we have presented high resolution \\halpha\\ imaging observations of Tycho, revealing the existence of a thin precursor which we interpret as a cosmic ray precursor.  While the current observation is consistent with inefficient acceleration, the observation of the precursor itself provides an important opportunity to constrain the key parameters of the acceleration, such as the diffusion coefficient and the injection parameters. A comparison with detailed numerical simulations will be critical to study the detailed physics of cosmic ray acceleration."
    },
    "1005/1005.3405_arXiv.txt": {
        "abstract": "{{\\it Context:} The expansion of network magnetic fields with height is a fundamental property of flux tube models. A rapid expansion is required to form a magnetic canopy.\\\\ {\\it Aims:} We characterize the observed expansion properties of magnetic network elements and compare them with the thin flux tube and sheet approximations, as well as with magnetoconvection simulations. \\\\ {\\it Methods:} We used data from the Hinode SOT NFI NaD$_{1}$ channel and spectropolarimeter to study the appearance of magnetic flux concentrations seen in circular polarization as a function of position on the solar disk. We compared the observations with synthetic observables from models based on the thin flux tube approximation and magnetoconvection simulations with two different upper boundary conditions for the magnetic field (potential and vertical). \\\\ {\\it Results:} The observed circular polarization signal of magnetic flux concentrations changes from unipolar at disk center to bipolar near the limb, which implies an expanding magnetic field. The observed expansion agrees with expansion properties derived from the thin flux sheet and tube approximations. Magnetoconvection simulations with a potential field as the upper boundary condition for the magnetic field also produce bipolar features near the limb while a simulation with a vertical field boundary condition does not. \\\\",
        "introduction": "At present, observations do not provide an explicit picture of how the chromospheric network magnetic field is structured. On one hand, we have increasing observational evidence of something which can be interpreted loosely as a canopy structure: e.g., fibrils in the Ca II infrared triplet lines; \\citealt{Vecchio+others2007}, large-scale canopy structures in combined Zeeman and Hanle studies: \\cite{Bianda+others1998, Stenflo+others2002}, canopy-like expansion seen in magnetograms near the limb; \\cite{Jones+Giovanelli1983}). On the other hand, \\add{we know that the Sun is more complicated than implied by simple models \\citep{WedemeyerBoehm09}. An important question is to what extent simple models of flux tubes are able to reproduce the center-to-limb appearance of the network magnetic field structures.}   The appearance of network magnetic flux concentrations in circular polarization maps changes from unipolar at disk center to bipolar near the limb. This is consistent with network magnetic fields expanding and fanning out with height as proposed by \\cite{Gabriel1976}. In the chromosphere, the fanning is clearly present; e.g., \\cite{Jones+Giovanelli1983} found low-lying, 200-800 km, magnetic canopies in magnetograms taken near the solar limb. More recently \\cite{Kontar+Hannah+MacKinnon2008} has used hard X-ray observations from RHESSI to estimate the expansion and found that the magnetic field expanded noticeably at a height of $\\approx$ 900 km.   Expansion of magnetic field with height has been studied in photospheric structures mostly using magnetic flux tube models. Indeed, flux tube models predict a rapid expansion of the field with height. \\cite{Solanki+others1999} shows that magnetic structures as different in size and flux as small flux tubes and sunspots have similar relative expansion rates, which agree with the thin flux tube approximation. A study of the characteristics of magnetic flux structures in radiative magnetohydrodynamic (MHD) simulations revealed the expansion properties to be similar with the thin flux tube and sheet approximations (\\citealt{Yelles+others2009}). The expansion is seen in observations: e.g., a thin flux tube model can simultaneously reproduce the observed Zeeman splittings of Mg 12.32 $\\mu$m, Fe 525.0 nm and Fe 1.56 $\\mu$m lines, which span the upper to the lower photosphere in formation height \\citep{Bruls+Solanki1995}. Bruls \\& Solanki also showed that a flux tube model can explain the Mg 12.32 $\\mu$m line profile shapes observed by Zirin \\& Popp (1989). Additional evidence of expansion is that a canopy resulting from the expansion of a flux tube can best explain the observed photospheric asymmetric Stokes $V$ profiles with weak zero-crossing shifts \\citep{Grossmann-Doerth+others1988}.  \\nocite{Zirin+Popp1989}  In this paper we use circular polarization maps from the Solar Optical Telescope (SOT, \\citealt{Tsuneta+others2008}) on the Hinode satellite to study the \\add{average} expansion properties by characterizing how the appearance of network flux concentrations changes from the solar disk center to the limb. The center-to-limb approach lets us examine the expansion at different viewing angles and at different heights due to the shift in the height \\add{range where} spectral lines \\add{are formed} as a function of $\\mu$ ($\\mu=cos(\\theta)$, where $\\theta$ is the viewing angle). To further expand the coverage we use SOT observations from the spectropolarimeter (SP) and the narrowband filter imager (NFI) NaD$_{1}$ channel. We combine the observations with modeling the expansion of magnetic flux with height \\add{at various $\\mu$-values} by using the thin flux tube and sheet approximations and more realistic 3-dimensional magneto-convection simulations. ",
        "conclusions": "The SOT observations show clearly the expansion of magnetic flux concentrations with height as a change from unipolar radial cuts at disk center to bipolar near the limb. Using the ratio of the radial cuts' center- to the limb-side peak as a proxy for expansion indicates that the magnetic field has fanned out more at the formation height of the NaD$_{1}$ NFI signal than at the formation height of the Fe I 630 nm SP signal. The NaD$_{1}$ ratios are smaller than the Fe I and the switch over from unipolar to bipolar cuts takes place closer to disk center, i.e., at $\\mu$=0.5 in NaD$_{1}$ rather than at $\\mu \\approx$0.3 in Fe I. Furthermore, unlike the SP data the NFI data exhibit exhibit only very few unipolar features close to the solar limb. The apparent expansion between the SP and NFI data and the lack of unipolar features close to the limb in NFI data can be caused by several factors: the NaD$_{1}$ signal is formed higher and the features that would appear unipolar in SP look bipolar in NFI due to increased expansion of the field with height, as expected from flux tube models. Alternatively, the unipolar features are not strong enough or do not reach high enough to be visible in the NFI data. The line formation of the NaD$_{1}$ and Fe I 630 nm lines are somewhat different, e.g., the effective $g$ value of the NaD$_{1}$ line is smaller than for either Fe line making it less sensitive to low magnetic field strengths or fluxes. The difference may be accentuated by other factors such as the NFI data being filtergram while the Fe I signals are deduced from spectrograph data. The different noise levels and pixel sizes of the instruments may also play a role. To fully compare the two one should use simultaneous NFI and SP observations at various disk positions.  Expansion based on the observed ratios and ratios derived from the thin flux sheet and tube approximations are in good agreement. The synthetic ratios are of similar magnitude and the switch over from unipolar to bipolar takes place at similar $\\mu$-values, i.e., around 0.5. Since the sheets expand faster, the thin sheet model results in bipolar features closer to disk center, first measurable ratios appear at $\\mu\\approx$0.8, and at smaller $r_{0}$: 200 km for the sheet and 400 km for the tube models (assuming the observations are perpendicular to the elongated direction of the sheet).   Unlike in the observations, in the thin sheet/tube models the ratios between the limb-to-center side Stokes $V$ amplitudes and areas of the Fe lines are smaller than the NaD$_{1}$ ratios. While the NaD$_{1}$ line core is formed higher than the Fe line cores, the line flanks, from which the bulk of the Stokes $V$ signal is from, are formed at roughly the same height in the sheet and tube models (based on response functions computed in LTE). The difference between the synthetic and observed ratios is probably at least partly due to the radiative transfer being done in LTE and a sampling bias due to the differences in the SP and NFI data. Also, factors such as model magnetic field strength and microturbulence may influence the ratios. If the disk center-side of the flux concentration has a strong enough field it may become Zeeman saturated before the limb-side, leading to a relatively reduced center side amplitude and a decreased ratio. Since the Fe lines have a larger Zeeman sensitivity this effect would first take place in the Fe lines.   The observations clearly show the expansion of the field which is found to be similar to the expansion derived from using the thin flux sheet and tube approximations. This result is in agreement with \\cite{Solanki+others1999}. In other aspects the magnetic network in the Sun is probably not well described by a simple flux tube or sheet scenario. In reality the network is dynamic and likely composed of individual flux concentrations which merge at some height. Observations of chromospheric lines, such as H-$\\alpha$ (e.g., \\citealt{DePontieu+others2007}), the Ca K (e.g., \\citealt{Zirin1974}, \\citealt{Pietarila+others2009}) and Ca II infrared triplet lines around 850 nm (e.g., \\citealt{Vecchio+others2007}), exhibit fibril structures which probably outline magnetic field lines that are more heated than others, i.e., the fibrillar canopy is thermally inhomogeneous. The magnetic field in contrast is likely quite smooth as indicated by the homogeneous appearance of both the uni- and bipolar patches in the SOT data.   The choice of upper BC for the magnetic field in the MHD-simulations has a significant effect on the expansion properties of the field. The BC affects not only the upper portion of the simulation domain but effects are also seen at the $\\tau$=1 level, namely the existence of horizontal and vertical field also outside the initial magnetic strip. In the potential field BC simulations horizontal magnetic field is present everywhere and a canopy-like structure is formed at the top of the domain whereas in the vertical field BC simulation the field remains confined in the strip. \\add{We plan to address in a later paper the time evolution of the magnetic flux in the potential field BC simulation to study in detail the transport of magnetic field from the top of the domain, i.e., the canopy, to the bottom of the photosphere.} \\add{Note that with the term canopy we are considering the canopy-like expansion of the field, not the actual chromospheric canopy.} The choice of BC affects strongly the synthetic radial cuts: no apparent bipolar features are seen in radial cuts from the vertical BC simulation. In contrast, a bipolar feature begins to appear at $\\mu=0.34$ in the potential BC case. The switch over from unipolar to bipolar takes place closer to the limb and the resulting ratios are larger than those seen in the observations or in the zeroth order thin flux tube/sheet models, indicating that the field does not fan out enough.   A stronger expansion in the potential field BC simulation could be achieved by, e.g., having two strips of opposite magnetic polarities in the initial setup or increasing the simulation domain size horizontally. The former would result in a canopy structure connecting the opposite polarities. However, based on the NFI ratios, the magnetic field in the mostly unipolar polar regions expands at a similar rate as in the equatorial limb regions, where mixed polarity fields are found. The latter alternative, larger simulation domain, would result in stronger fanning out because of the periodic BCs in the horizontal direction. Currently the domain size is 24 Mm which is comparable to the typical size of internetwork cells (30 Mm, \\citealt{Beckers1968}). Therefore a larger domain would not necessarily be more realistic. Placing the upper boundary higher might also result in stronger expansion \\add{since the field would have more volume to expand in}. However, the current treatment of radiative transfer in LTE would become increasingly incorrect. Finally, the field may not yet have reached a potential state at the upper boundary, so that even this boundary condition may not be appropriate.  We plan to extend the current study by using spectropolarimetric observations of photospheric and chromospheric lines to study in detail the network magnetic fields.  \\add{ \\subsection{Conclusions} The circular polarization signal of magnetic flux concentrations is known to change from unipolar at disk center to bipolar near the limb, consistent with the magnetic field fanning out with height. \\cite{Solanki+others1999} showed with observations then available that the relative expansion rates of features of various sizes, were consistent with rates derived for simple thin flux tubes. The data from the Hinode satellite, with their high spatial resolution, low noise and freedom from atmospheric seeing effects, provide a new opportunity to address the question with better statistics and with an emphasis on the variation from disk center to the limb. We again found the observations to be in good agreement with expansion properties derived from the thin flux sheet and tube approximations, confirming the previous results with the new data and a more extensive analysis. Since the signal we are interested in forms higher in the atmosphere near the limb, we found that realistic numerical simulations, which extend only to 700 km above $log(\\tau_{500nm})=0$, are less useful in modeling these observations, with the results strongly depending on the choice of boundary conditions for the magnetic field. }"
    },
    "1005/1005.2709_arXiv.txt": {
        "abstract": "We present wide-field imaging of the 2007 outburst of Comet 17P/Holmes obtained serendipitously by SuperWASP-North on 17 nights over a 42-night period beginning on the night (2007 October 22-23) immediately prior to the outburst.  Photometry of 17P's unresolved coma in SuperWASP data taken on the first night of the outburst is consistent with exponential brightening, suggesting that the rapid increase in the scattering cross-section of the coma could be largely due to the progressive fragmentation of ejected material produced on a very short timescale at the time of the initial outburst, with fragmentation timescales decreasing from $t_{frag}\\sim2\\times10^3$~s to $t_{frag}\\sim1\\times10^3$~s over our observing period.  Analysis of the expansion of 17P's coma reveals a velocity gradient suggesting that the outer coma was dominated by material ejected in an instantaneous, explosive manner.  We find an expansion velocity at the edge of the dust coma of $v_{exp}=0.55\\pm0.02$~km~s$^{-1}$ and a likely outburst date of $t_0=2007~{\\rm October}~23.3\\pm0.3$, consistent with our finding that the comet remained below SuperWASP's detection limit of $m_V\\sim15$~mag until at least 2007 October 23.3.  Modelling of 17P's gas coma indicates that its outer edge, which was observed to extend past the outer dust coma, is best explained with a single pulse of gas production, consistent with our conclusions concerning the production of the outer dust coma. ",
        "introduction": "Discovered on 1892 November 9 by Edwin Holmes, Comet 17P/Holmes (hereafter, 17P) is a dynamically ordinary Jupiter-family comet, with a semimajor axis of $a=3.617$~AU, an eccentricity of $e=0.432$, an inclination of $i=19.11^{\\circ}$, an orbital period of $P_{orb}=6.8$~years, and a Tisserand parameter (with respect to Jupiter) of $T_J=2.859$.  On 2007 October 24.1, it was discovered by J. A. Henr\\'iquez Santana to be undergoing a substantial outburst.  Starting from an initially observed magnitude of $m_V\\sim8.4$ \\citep[compared to a pre-outburst brightness of $m_V\\sim17$ measured on 2007 October 23.1;][]{cas07}, it brightened by approximately 0.5 mag~hr$^{-1}$ over the course of 6 hours \\citep{buz07}.  The comet reached $m_V\\sim2.0$ by 2007 October 25.1 \\citep{spo07,elh10}, representing a million-fold increase in brightness over only 2 days.  This outburst of 17P mirrored a similar outburst that led to the comet's initial discovery in 1892, and continued to be monitored by both amateur and professional astronomers as it progressed.  Spectroscopic analyses indicated that the coma was largely dominated by dust grains \\citep{sch07}, but evidence of cometary volatiles and sublimation byproducts, including OH, H$_2$O, NH, CN, HCN, HNC, C$_2$, C$_3$, C$_2$H$_2$, C$_2$H$_6$, CS, CH$_3$CN, and CH$_3$OH, were also detected during the outburst \\citep{dra07,fit07,kob07,sch07,wag07,boc08,del08,yan09}. Negative polarisation was found for the dust coma by \\citet{jos10}, consistent with measurements made for other comets. A delay in the appearance of CN, C$_2$, and [OI] emission lines between 2007 October 24.58 and October 25.46 led \\citet{kob07} to speculate that their observations could be at least partially explained by the delayed sublimation of icy grains ejected at the time of the outburst.  This hypothesis is supported by \\citet{del08} who found evidence of a distributed source of sublimation products, consistent with the ejection and subsequent delayed sublimation of icy grains in the coma. The total mass loss due to this outburst has been estimated to be $\\sim10^{12}$ kg or a few percent of the comet's total mass \\citep{sek08a,mon08}. ",
        "conclusions": "\\subsection{Primary Findings}  Wide-field imaging of the 2007 outburst of 17P/Holmes was serendipitously obtained by the SuperWASP-North facility on 17 nights over a 42-night period beginning on the night (2007 October 22-23) immediately prior to the outburst.  We report the following key findings: \\begin{itemize} \\item{The comet was not detected in data from the night before the outburst's discovery on the night of 2007 October 22-23, indicating that it remained below the SuperWASP detection threshold of $m_V\\sim15$~mag (consistent with pre-outburst reports placing the comet's brightness at $m_V\\sim17$~mag) until at least 2007 October 23.3. } \\item{The unresolved coma (as seen by SuperWASP) was likely optically thin during our observations on the first night of the outburst.  The comet's lightcurve during these observations is consistent with an exponential function, suggesting that its rapid brightening could have been driven by the progressive fragmentation of ejected material produced on a very short timescale at the time of the initial outburst. Our best-fit functions to the data imply an initial fragmentation timescale of $t_{frag}\\sim2\\times10^3$~s, decreasing to $t_{frag}\\sim1\\times10^3$~s, and a near-instantaneous leap in brightness from a magnitude of $m_V\\sim17$~mag to $m_V\\sim11$~mag at the moment of the initial outburst. } \\item{Analysis of coma surface brightness profiles reveals a velocity gradient consistent with the outer coma being dominated by material ejected in an explosive manner (i.e., at a single instant), rather than over an extended period of time.  Near the outer edge of the visible coma, where the coma likely transitions from being dust-dominated to gas-dominated, this velocity gradient corresponds to an expansion velocity of $0.55\\pm0.02$~km~s$^{-1}$, consistent with previous reported measurements.  From this analysis, we find a most likely outburst date of $t_0=2007~{\\rm October}~23.3\\pm0.3$. This finding of explosively ejected material dominating the outer coma does not rule out the possibility of subsequent sustained mass ejection supplying the inner coma at later times. } \\item{We measure the rate of motion (relative to the nucleus) of the secondary brightness peak in 17P's coma, which we refer to as the ``false nucleus'', and find a projected drift velocity of $v_{drift}=120\\pm5$~m~s$^{-1}$, consistent with previously reported measurements. } \\item{Dust modelling shows that 17P's nucleus is a plausible primary emission source of outer coma material, and that a secondary source such as a separated nucleus fragment is not required to explain the motion of the coma relative to the nucleus. We show instead that the drifting of the coma relative to the nucleus can be explained as the consequence of radiation pressure alone. } \\item{Modelling of 17P's gas coma indicates that the morphology of the portion of the observed coma profile hypothesised to be gas extending past the outer dust coma is best explained by a single, instantaneous outburst of gas production, rather than extended gas production that persists over several days. This result is consistent with our conclusion that the outer dust coma is likely dominated by material ejected instantaneously, and not over an extended period of time. We also note that C$_2$ is likely to be the dominant observed component of the gas coma, and find that its decay time is consistent with SuperWASP observations of the disappearance of the observed coma profile region attributed to gas. } \\end{itemize}    \\subsection{Anatomy of an Outburst}  To date, no consensus has emerged to explain the physical origin of 17P's spectacular outbursts in 1892 and 2007.  A primary difficulty lies in reconciling the occurrence of such an apparently catastrophic event at least twice in the recent past with a lack of comparable outbursts in almost all other comets, although \\citet{sek08b} has noted that 1P/Halley exhibited a similarly powerful outburst in 1836. \\citet{whi84} suggested that a grazing encounter and eventual impact by a small satellite with the nucleus could have been responsible for the 1892 outburst, but such an explanation was rendered highly implausible with a second episode of outburst activity in 2007. Furthermore, any model formulated to explain 17P's behaviour must also account for the lack of outburst activity in the 115 years between 1892 and 2007 \\citep{sek09b}.  \\citet{sek08a,sek08b} has proposed a scenario in which an inwardly-diffusing thermal wave gradually penetrated a large, weakly-cemented pancake-shaped layer of the nucleus \\citep[analogous to those observed for 9P/Tempel 1; {\\it cf}.][]{tho07} over numerous orbits until it finally reached a large reservoir of amorphous water ice at the layer's base. Upon being heated to the necessary transition temperature, this amorphous ice layer underwent an exothermic transformation to crystalline ice, leading to rapid sublimation and causing the pancake-shaped layer to separate from the nucleus and almost immediately explosively disintegrate.  This two-part outburst scenario is invoked by Sekanina to explain the appearance of the false nucleus as the central source of 17P's dust coma, as well as its apparent motion relative to the true nucleus.  We note, however, that our dust modelling (\\S\\ref{dustmodel}) shows that it is not necessary for the false nucleus to be the source of emission to explain the position of the outer dust coma relative to the true nucleus.  Thus a model where a large surface fragment first violently separates from the nucleus (with a projected velocity exceeding 100~m~s$^{-1}$) and then explosively disintegrates, requiring two non-simultaneous catastrophic events, may be needlessly complicated.  The false nucleus could still be a concentration of material related to a large fragment ejected in the initial outburst that then disintegrated, but may be better characterised as a byproduct, rather than a key component, of the outburst.  \\citet{alt09} propose an alternate hypothesis in which a particularly close perihelion passage by 17P on 2007 May 4 substantially raised the sublimation rate of subsurface water ice relative to previous perihelion passages.  This increased production of subsurface water vapour, coupled with an ``airtight'' layer of surface regolith, ultimately led to an explosive disintegration of 17P's dust mantle 172 days later, causing the observed outburst. While it is true that 17P's perihelion passage just 2.05~AU from the Sun was its closest approach to the Sun in over a century (according to JPL's Horizons ephemeris generator at http://ssd.jpl.nasa.gov/horizons.cgi), we note that 17P's perihelion distance immediately prior to its 1892 outburst was a more modest 2.14~AU from the Sun, a distance comparable to perihelion passage distances in 1864, 1871, 1878, 1885, 1899, 1906, 1972, 1979, 1986, 1993, and 2000, all of which were reached without any recorded reports of behaviour comparable to the comet's 1892 or 2007 outbursts.  Most recently, \\citet{rea10} propose a scenario in which trapped gases released by the crystallisation of amorphous ice, as well as sublimation of other ices driven by the exothermic crystallisation process, cause a buildup of gas pressure in a subsurface cavity on the comet.  The outburst of 17P then occurred when the pressure built up within this cavity exceeded the strength of the surrounding material (found to be relatively high --- $10-100$~kPa --- based on the nucleus's survival of the outburst), causing the cavity to rupture violently and energetically, creating the observed explosive ejection of nuclear material. We note that while \\citet{kos10} report that the crystallisation of amorphous water ice is unlikely to have caused 17P's outburst, they reach this conclusion based on a simple model in which the ice being crystallised is present in a sub-surface layer just below a dust mantle.  The hypothesis presented by Reach {\\it et al.} is based on a different physical scenario where ice crystallisation occurs in a subsurface void in which gas pressure is able to build up, and as such, we believe that it remains plausible and, of the hypotheses proposed thus far, is the most consistent with our own findings. We note, however, that it then raises questions as to why 17P is the only comet known thus far with the apparently unique combination of high-tensile-strength nuclear material, subsurface cavities, and amorphous ice capable of driving two of the largest cometary outbursts observed in modern astronomy.  \\citet{gai07} reported that observations they made between 2007 October 24 to 2007 November 4 with the Pic du Midi 1~m telescope show multiple dust streams with well-defined origin points, four of which were measured to recede from 17P's nucleus at roughly constant velocities (as projected on the sky) ranging from 50 to 100~m~s$^{-1}$, implying the presence of unresolved, steadily disintegrating nucleus fragments in the coma. These fragments were also calculated to have separated from the nucleus between 2007 October 23.7 and 2007 October 24.8. Additionally, \\citet{ste09} observed at least sixteen 10-m- to 100-m-scale fragments receding from the nucleus that showed evidence of ongoing sublimation and disintegration. Coupled with our finding of dust fragmentation in the first hours after outburst, these observations point to a scenario of continued fragmentation of cometary material following the outburst.  We certainly still have far from a complete picture of 17P's 2007 outburst. We suggest that detailed dust modelling and analysis of images of the comet's inner coma using data with higher spatial resolution than SuperWASP would be useful for clarifying the temporal and kinematic nature of the outburst, in particular whether significant dust production continued from the nucleus after the initial outburst event and, if so, whether ejection velocities were comparable to those in the initial outburst event. More detailed discussions of spectroscopy would also help constrain the spatial and temporal nature of the contribution of gaseous species to the 17P's appearance, such as whether sublimation of nucleus-bound ices or ice particles in the coma was significant, and whether there was any appreciable delay in the sublimation of icy material in the coma that could be linked to hierarchical fragmentation of macroscopic nucleus particles."
    },
    "1005/1005.5259_arXiv.txt": {
        "abstract": "This paper considers magnetic field generation by a fluid flow in a system referred to as the Archontis dynamo: a steady nonlinear magnetohydrodynamic (MHD) state is driven by a prescribed body force. The field and flow become almost equal and dissipation is concentrated in cigar-like structures centred on straight-line separatrices. Numerical scaling laws for energy and dissipation are given that extend previous calculations to smaller diffusivities. The symmetries of the dynamo are set out, together with their implications for the structure of field and flow along the separatrices. The scaling of the cigar-like dissipative regions, as the square root of the diffusivities, is explained by approximations near the separatrices. Rigorous results on the existence and smoothness of solutions to the steady, forced MHD equations are given. ",
        "introduction": "Much is known about fast dynamo action: the rapid growth of magnetic fields at high magnetic Reynolds number in fluid flows with chaotic streamlines, but the mechanisms for the dynamical saturation of such fields remain poorly understood. In many cases when the growing field equilibrates by modifying the fluid motion, the effect is to switch off the chaotic stretching in the flow, as measured for example by a reduction in the finite-time Liapunov exponents \\cite[e.g.,][]{CaHuKi96,ZiPoPo98}. What is left is a fluid threaded by a magnetic field which resists stretching and so suppresses overturning fluid motions, but supports elastic wave-like motions, essentially Alfv\\'en waves with coupled field and flow \\cite[e.g.,][]{CoHuPr10}. The final state of many simulations shows apparently chaotic behaviour in space and time, suggestive of an attractor of moderate or high dimension, although because of the three-dimensionality of MHD systems little can be done to explore its properties, for example the fractal dimension or spectrum of Liapunov exponents.  Although this appears to be the outcome of many simulations, as far as they can be run, there are some intriguing examples where a further phase of evolution takes place: the magnetic field and flow align, depleting the nonlinear terms, and both fields evolve to a steady (or very slowly evolving) state. The key point is that in unforced, ideal magnetohydrodynamics (see equations (\\ref{eqNS0}--\\ref{eqdivfree0}) below with $\\nu=\\eta=0$ and $\\fv=0$) any state with $\\uv = \\pm \\bv$ is an exact steady solution. The remarkable fact that simulations of forced, non-ideal MHD turbulence could evolve to something very close to such a state was first observed by \\cite{Ar00} in his thesis, and published in \\cite{DoAr04} (hence referred to as DA), and \\cite{ArDoNo07}. These simulations use a compressible code with a Kolmogorov forcing function, (\\ref{eqSvdef}) below, first used as the form of a flow for simulations of fast, kinematic dynamo action by \\cite{GaPr92}. Subsequently \\cite{CaGa06a} undertook incompressible simulations of the same system as Archontis, and pushed up the fluid and magnetic Reynolds numbers; our work is linked closely to this paper, which we refer to as CG in what follows.  What these authors found was that, starting with a forced fluid flow and a seed magnetic field, the growing magnetic field initially equilibrates in rough equipartition with the velocity field, in a messy, chaotic time-dependent state. However during this state, there is a slow but persistent exponential growth in the average alignment of the $\\uv$ and $\\bv$ vectors, as measured by the cross-helicity. This process of alignment continues until there takes place a sudden increase in the fluid and magnetic energies, and both fields tend to a steady state of almost perfect alignment, discrepancies being controlled by the weak dissipation and the forcing. In fact since any solution $\\uv=\\pm\\bv$ is a neutrally stable solution of the ideal problem \\citep{FrVi95}, the solution that is selected must depend delicately on balances involving these subdominant diffusive and forcing effects. We note that some alignment of field and flow has been noted in many other MHD flows, for example see \\cite{DoMaVe80}, \\cite{PoMeFr86}, \\cite{MaCaBo06} and references therein, but of a less dramatic nature.  This observation of dynamo saturation in a steady state with such a high degree of alignment was a new phenomenon: CG refer to the saturated state as the `Archontis dynamo', though we prefer the term `Archontis saturation mechanism'. CG observed this aligned  state as a solution branch over a wide range of magnetic and fluid Reynolds numbers (taking the magnetic Prandtl number to be unity in much of their work). Further developments include the development of bursts of rapid time dependence after some time in the steady state, in the study \\cite{ArDoNo07}. However this appears only to occur in the compressible case, as it has not been seen by CG nor in our simulations; we will therefore not discuss this further. \\cite{CaGa06b} also find slow time-dependent evolution of the saturated state for the Kolmogorov forcing with magnetic Prandtl number $\\Pra = \\nu/\\eta$ not equal to unity, and for more general spatially periodic steady forcings. In all cases though, the field and flow settle into a state of very close alignment, even if they then evolve on a slow time scale.  The focus of the present paper is to understand more about the structure of the steady saturated state for the Kolmogorov forcing and unit magnetic Prandtl number $\\Pra$, with a particular focus on the regions where dissipation occurs and on rigorous results on existence and smoothness. DA and CG find a complex geometrical picture for the field and flow and identify these regions of high dissipation: they are localised along straight-line separatrices that join a family of stagnation points; similar structures are found in the 1:1:1 ABC flow \\citep{DoFrGrHeMeSo86}. These are found to have a width scaling as $\\sqrt{\\eps}$ where $\\eps$ is a dimensionless measure of the diffusivity, and one of our aims is to understand this power law.  We set up the governing equations in \\S\\ref{secgov} and extend the solution branch to yet smaller values of the diffusivity $\\eps$ by means of large scale simulations in \\S\\ref{secnum}. In \\S\\ref{secsym} we then classify the symmetries of the Kolmogorov forcing, which are preserved by the nonlinear, saturated field and flow. These symmetries are the reason for the presence of the non-generic straight line separatrices that join stagnation points in the flow and field, and they constrain the local flow: it is in these regions that dissipation is strongest. We plot the local structure of fields along the separatrix from $(0,0,0)$ to $(\\pi,\\pi,\\pi)$ in \\S\\ref{secflowfield}. We determine the effects of diffusion by setting up PDEs for the advection of field as it enters the dissipative regions in \\S\\ref{secPDE} and use these to justify the order $\\sqrt{\\eps}$ scaling for the cigar widths found in CG. We then proceed with a formal mathematical investigation of the existence of steady-state solutions to the MHD problem at hand and bounds for them in various function spaces in \\S\\S \\ref{maths}--\\ref{smooth}. The reader should note that these sections use functional analysis and so have a different flavour from the earlier ones. Finally \\S\\ref{secdisc} offers concluding discussion. ",
        "conclusions": "\\label{secdisc}  We have presented investigations into the structure of the magnetic field and flow in the equilibrated regime of the Archontis dynamo. Because of the highly three-dimensional nature of the system, application of the available analytical tools yields only rough results of limited value, and we lack any kind of complete solution. What we have done is first to extend the range of diffusivities $\\eps$ over which the saturation mechanism operates to give the steady state with nearly aligned fields. We have also classified the symmetries of these flows and measured the field structure on the separatrices, home of the cigar-like dissipative regions.  Then, using basic analytical tools, we have investigated the scaling of diffusive terms near the separatrices. Here at leading order the field $\\lamv=\\eps^{-1}\\Lamv_-$ that enters from the body of the flow is transported along characteristics of $\\Lamv = \\Lamv_+$. Where these characteristics come together in the compressive flow at the stagnation points, where trajectories spiral in, large gradients in $\\lamv$ are generated, and diffusive terms enter the problem on scales of $\\eps^{1/2}$ as found by CG. In more general flows we may expect a similar behaviour, with regions of heightened dissipation localised at points where $\\Lamv=0$ and along the unstable manifolds of such points. Of course in the Archontis example the unstable manifolds link the stagnation points and so the topology here is very simple and the dissipative regions very small, of order $O(\\eps)$ in volume: in other cases they may wander through the three-dimensional space, giving a picture of much greater complexity, as could be occurring in examples in \\cite{CaGa06b}. Again wider regions of dissipation, perhaps dense in the space, could occur if examples exist where $\\Lamv$ has no stagnation points; unfortunately the form of $\\Lamv$ is not under our control except where strongly constrained by symmetries. In order to cope with unknown levels of geometrical complexity, an approach based on functional analysis is appropriate, and this is the final part of the paper, in which the existence and smoothness properties of steady solutions are established.  An analogy of the naively truncated equations (namely (\\ref{eqLamv}, \\ref{eqlamv}) with $\\eps=0$) with the Euler equation, which is the subject of intense research, is instructive. A method for the investigation of the evolutionary Euler and Navier--Stokes equations consists of the introduction into the equations of new regularising terms, such as $\\eps(-\\nabla^2)^\\alpha\\uv$ or $\\eps(-\\nabla^2)^\\beta(\\partial\\uv/\\partial t)$. It has been known for decades that for $\\alpha>5/4$ solutions to the regularised equations are infinitely differentiable at any $t>0$; for $\\beta\\geq1 /2$ and $\\beta > 5/6$ one can prove analyticity, at any $t>0$, of solutions to the Navier--Stokes and Euler equations, respectively \\citep{Zh10}. For any $\\alpha$ or $\\beta$ below the respective thresholds, the problem is as difficult as the one for the original equation. When the limit $\\eps\\to0$ is considered, the results are so far inconclusive. One can only show that there exist sequences $\\eps_k\\to0$ such that solutions for these $\\eps_k$ converge to a weak solution to the non-regularised equation, and either the limit weak solution is unique for all such sequences, or there exists a continuum of weak solutions. Whether for $\\eps\\to0$ singularities develop in derivatives of the regularised solutions, and how strong they are if they develop, remains unknown.  The difficulties arise in the general theory, because the bounds for solutions are singular in $\\eps$ as $\\eps\\to0$. Here the analogy with the Archontis dynamo problem crystallises: in the Archontis problem the diffusive terms can be regarded as a regularisation of the naively truncated diffusionless problem, and we need to find out what happens when the regularisation parameter $\\eps$ tends to zero. (In the diffusionless, i.e.\\ non-regularised, case it is unclear whether weak steady solutions exist.) We note that the analogy may work both ways: the asymptotic analysis near the separatrix in the Archontis dynamo (which we present in \\S\\S \\ref{secflowfield} and \\ref{secPDE}) may contain clues to what happens in solutions to the regularised Euler (or even Navier--Stokes) equations in the limit $\\eps\\to0$. Unfortunately, the clues are well hidden, because the regularising term in the Archontis problem is of a different structure, and a very specific symmetric steady solution to the general system of MHD equations is considered.  Besides further attempts to carry out an asymptotic analysis of equations \\rf{eqLamv} and \\rf{eqlamv} and their evolutionary versions, a number of other directions could be pursued in the future, for example investigating time-dependent modifications to the steady Kolmogorov forcing used here, or studying the evolution of superposed large-scale fields and corresponding non-helical transport effects, as in the recent work of \\cite{SuBr09}."
    },
    "1005/1005.3155_arXiv.txt": {
        "abstract": "We investigate the far-infrared--radio correlation (FRC) of stellar-mass-selected galaxies in the Extended Chandra Deep Field South using far-infrared imaging from \\textit{Spitzer} and radio imaging from the {Very Large Array} and {Giant Metre-Wave Radio Telescope}. We stack in redshift bins to probe galaxies below the noise and confusion limits. Radio fluxes are $K$--corrected using observed flux ratios, leading to tentative evidence for an evolution in spectral index. We compare spectral energy distribution (SED) templates of local galaxies for $K$--correcting FIR fluxes, and show that the data are best fit by a quiescent spiral template (M51) rather than a warm starburst (M82) or ULIRG (Arp220), implying a predominance of cold dust in massive galaxies at high redshift.  In contrast we measure total infrared luminosities that are consistent with high star-formation rates. We observe that the FRC index ($q$) does not evolve significantly over $z=0-2$ when computed from $K$-corrected 24 or 160-\\mum\\ photometry, but that using 70-\\mum\\ fluxes leads to an apparent decline in $q$ beyond $z\\sim1$.  This suggests some change in the SED at high redshift, either a steepening of the spectrum at rest-frame $\\sim25-35$\\mum\\ or a deficiency at $\\sim70$\\mum\\ leading to a drop in the total infrared/radio ratios.  We compare our results to other work in the literature and find synergies with recent findings {on the high-redshift FRC, high specific star-formation rates of massive galaxies and the cold dust temperatures in these galaxies}. ",
        "introduction": "One of the most exciting research areas in observational astronomy at this time is the field of far-infrared astronomy.  Our understanding of extragalactic sources in the far-infrared (FIR) and sub-millimetre (sub-mm) to millimetre regimes has improved exceptionally over the past decade, thanks to such instruments as ESA's \\textit{Infrared Space Observatory} (\\textit{ISO}, launched 1995; \\citealp{kessler_infrared_1996}) and NASA's \\textit{Spitzer Space Telescope} (launched 2003; \\citealp{werner_spitzer_2004}), alongside ground-based instruments such as the Sub-millimetre Common-User Bolometer Array (SCUBA; \\citealp{holland_scuba:common-user_1999}) and the Max Planck Millimeter Bolometer Array (MAMBO; \\citealp{kreysa_bolometer_1998}), both commissioned in the late 1990's.  One of the most important FIR instruments to date is the Multiband Imaging Photometer for \\textit{Spitzer} \\citep[MIPS;][]{rieke_multiband_2004}, which, with \\textit{Spitzer}'s 0.85m mirror provides imaging with diffraction limited resolutions of 6, 18 and 40 arcsec in three broad bands centred at 24, 70 and 160\\mum\\ respectively.  Recent advances in the resolution of the cosmic infrared background (CIB) by stacking into images from \\textit{Spitzer} {\\citep[see e.g.][]{dole_cosmic_2006,dye_scuba/spitzer_2007,chary_new_2010}} and BLAST \\citep[The Balloon--borne Large Aperture Sub-mm Telescope;][]{marsden_blast:_2009} have enabled an improved understanding of the history of star formation and galaxy formation and evolution, and this will be further improved by ongoing work \\citep[such as][]{berta_dissecting_2010} with ESA's new \\textit{Herschel Space Observatory} \\citep{pilbratt_herschel_2010}.  Most of the stellar mass in the local universe is concentrated in the most massive galaxies ($M_\\star \\gtrsim 10^{11}\\text{M}_\\odot$; \\citealp{kauffmann_stellar_2003}) and observations show that these have been in place since $z\\sim 1$ \\citep{dickinson_evolution_2003, bundy_mass_2005, prez-gonzlez_stellar_2008, taylor_rise_2009, collins_early_2009}.  The significant increase in density of luminous ($L_{8-1000\\mum} \\gtrsim 10^{11} \\text{L}_\\odot$) and ultra-luminous ($L_{8-1000\\mum} \\gtrsim 10^{12} \\text{L}_\\odot$) infrared galaxies (LIRGS and ULIRGs) from the local universe to $z\\sim 2-3$ \\citep[e.g.][]{daddi_passively_2005, daddi_population_2005, caputi_role_2006} is thought to reveal the formation stages of these latter-day giants, which apparently formed in a remarkably short time between $1 \\lesssim z \\lesssim 3$, in an antithesis to the paradigm of hierarchical structure formation \\citep[e.g.][]{de_lucia_formation_2006}.  \\label{sec:sources} In the FIR--sub-mm regime the dominant source of the extragalactic background light (after the cosmic microwave background) is thermal continuum emission from interstellar dust, which is mainly composed of polycyclic aromatic hydrocarbons (PAHs), graphites and silicates, typically less than $\\sim 0.25\\mum$\\ in size \\citep{draine_infrared_2007, draine_dust_2007}.  The FIR emission in star-forming galaxies is thought to arise both from cold dust in the large-scale `cirrus' component of the interstellar medium (ISM), and from warmer dust in and around star-forming regions \\citep[e.g.][]{de_jong_iras_1984,helou_iras_1986}.  If there is a sufficient level of dust-enshrouded star formation then the FIR emission is dominated by this `warm' dust component, which has characteristic temperatures of around 30--50K \\citep{dunne_scuba_2001, sajina_1-1000m_2006, dye_scuba/spitzer_2007, pascale_blast:far-infrared_2009}.  The dust is heated to these temperatures by the ultraviolet (UV) radiation field from hot O and B type stars, which are present only in regions of ongoing or recent star formation (`recent' meaning within the lifetime of these short-lived stars, $\\lesssim10^{8}\\text{yr}$; \\citealp{kennicutt_star_1998}). For this reason, the total IR luminosity ($L_\\text{TIR} = L_{8-1000\\mum}$) can be used as a tracer of star-formation rates (SFRs) in galaxies, often using one or a number of FIR fluxes (such as the \\textit{Spitzer} MIPS bands) or sub-mm fluxes to estimate $L_\\text{TIR}$ \\citep{kennicutt_star_1998}.  One issue with using the FIR as a SFR tracer is the contribution from cold dust in the ISM, which is heated by older stars in the disk of the galaxy, and is therefore unrelated to star formation \\citep{calzetti_calibration_2010}. Contamination of samples by galaxies hosting active galactic nuclei (AGN) is another problem, as AGN also emit UV radiation which can heat dust in the torus.  AGN-heated dust is generally hotter than dust heated in star-forming regions, so the thermal spectrum peaks at a shorter wavelength, and mid-infrared (MIR) fluxes (including \\textit{Spitzer}'s 24-\\mum\\ band, as well as shorter wavelengths) would be boosted.  Mid-infrared fluxes are also affected more uncertainly by the emission features of PAH molecules, which are ubiquitous in star-forming galaxies \\citep[e.g.][]{leger_identification_1984, roche_atlas_1991, lutz_nature_1998, allamandola_modeling_1999}, and the 10-\\mum\\ silicate absorption trough, so the use of 24-\\mum\\ fluxes as SFR indicators at high redshifts is subject to some contention \\citep[see e.g.][]{dale_infrared_2005, calzetti_calibration_2007, daddi_multiwavelength_2007, papovich_spitzer_2007, young_mid-_2009, rieke_determining_2009}.  \\label{sec:FIRRC} Another part of a galaxy's SED that can be used as a SFR indicator is the radio luminosity.  Non-thermal radio continuum emission from star-forming galaxies originates from type II supernova remnants (SNRs), the endpoints of the same massive short-lived stars that heat the dust via their UV radiation.  This connection with dust heating is important because it leads to the well-known (but not fully understood) FIR--Radio correlation \\citep[FRC;][etc]{van_der_kruit_high-resolution_1973, rickard_far-infrared_1984, helou_thermal_1985, condon_radio_1992}. The FRC is linear, remarkably tight and holds for a wide range of galaxy types over at least five orders of magnitude in luminosity \\citep{yun_radio_2001}.  It can be explained in terms of ongoing star~formation producing hot massive ($M>8\\text{M}_\\odot$) stars: while the FIR flux is emitted from dust heated by these stars, the radio emission arises from synchrotron radiation by cosmic ray (CR) electrons accelerated in the SNRs of the dying stars.  The non-thermal radio emission is smeared out through the galaxy as the relativistic CR electrons travel through the galaxy over lifetimes of $\\tau \\sim 10^{8}$ years, during which they emit synchrotron radiation via interactions with the galactic magnetic field {(as described by \\citealt{condon_radio_1992}, and shown observationally by \\citealt{murphy_effect_2006})}.  A shallower thermal component is also present in the radio spectrum due to bremsstrahlung radiation from electrons in H\\,{\\sc ii} regions, but this becomes dominant only at high frequencies ($\\gtrsim 30$~GHz) and at 1.4~GHz only comprises $\\sim 10\\%$ of the radio flux \\citep{condon_radio_1992}.  It is difficult however to explain the linearity and tightness of the correlation between thermal FIR luminosity and non-thermal radio luminosity.  `Minimum energy' estimates of magnetic fields in galaxies \\citep{burbidge_synchrotron_1956}\\footnote{ Inferring magnetic field strengths of galaxies is problematic and depends on many unknowns. The minimum energy argument circumvents some of these problems by specifying the minimum of the total energy density (a function of field strength), which occurs at the point where the energy density in particles is in approximate equipartition with that in the magnetic field \\citep{longair_high_1994}. Although there is no physical requirement for this equipartition to occur, it does provide an order-of-magnitude estimate for the magnetic field energy density, and is physically motivated if one considers that the magnetic field is tangled by turbulent motions in the interstellar plasma, and that these motions efficiently accelerate cosmic ray particles (see \\citealt{longair_high_1994} and \\citealt{thompson_magnetic_2006}). Equipartition has been observationally confirmed locally in the Milky Way \\citep{strong_diffuse_2000} and in normal spiral galaxies \\citep{vallee_galactic_1995,beck_magnetic_2000}, although its validity in starbursts and ULIRGs is less certain \\citep{thompson_magnetic_2006}.} imply a large variation between normal galaxies and extreme starbursts like Arp220.  To explain a constant FIR/radio ratio between such disparate systems, complex physical solutions need to be provided, for example invoking strong fine-tuning to regulate electron escape and cooling timescales, or short cooling timescales with magnetic fields $\\sim10$ times stronger than the `minimum energy' argument suggests \\citep{thompson_magnetic_2006}.  \\Citet{voelk_correlation_1989} first suggested a `calorimeter' model whereby both UV light from massive stars and CR electrons from SNRs are (a) proportional to the supernova rate, and (b) efficiently absorbed and reprocessed within the galaxy, so that the respective energy outputs in FIR re-radiation and radio synchrotron would both be tied to the supernova rate.  This theory requires a correlation between the average energy density of the radiation field and the galaxy magnetic field energy density. \\Citet{voelk_correlation_1989} argues this is plausible if the origin of the magnetic field is a turbulent dynamo effect, since the turbulence would be largely caused by the activity of massive stars, and hence correlated with the supernova rate.  Alternative non-calorimetric models include those of \\citet{helou_physical_1993}, using a correlation between disk scale height and the escape scale length for CR electrons; and \\citet{niklas_new_1997}, in which the FRC is driven by correlations with the overall gas density and equipartition of magnetic field and CR energy.  \\Citet{bell_estimating_2003} argues for a `conspiracy' to diminish both the FIR and radio emission originating from star formation in low luminosity galaxies when compared with luminous $\\sim L_\\star$ galaxies -- without this the relationship would not remain linear over the full luminosity range. Similarly, the calorimeter model of \\citet{lacki_physics_2009-1} invokes conspiracies in low and high gas surface density regions to maintain the relationship.  The physical origin of the FRC therefore is still an open question.  A full review of the theories is beyond the scope of this study, but a more detailed discussion of the literature can be found in \\citet{vlahakis_far-infrared-radio_2007}, and a more in-depth treatment is provided {in the numerical work of} \\citet{lacki_physics_2009-1}.    An ongoing strand of research at the current time is the investigation of the FRC at high redshifts and low fluxes, in particular whether there is any evolution \\citep[e.g.][]{ appleton_far-_2004, frayer_spitzer_2006, ibar_exploringinfrared/radio_2008, garn_relationship_2009, seymour_investigatingfar-ir/radio_2009, ivison_blast:far-infrared/radio_2009,sargent_vla-cosmos_2010,sargent_no_evolution_2010}. Measurements of any evolution (or lack thereof) would improve the accuracy of FIR-/radio-estimated SFRs at high redshift, and could shed light on the mechanism governing the FRC, as well as highlighting differences in the physical and chemical properties of star-forming galaxies at high and low redshift \\citep{seymour_investigatingfar-ir/radio_2009}.   In the current work we investigate the FRC over a range of redshifts, for a sample that is not limited by FIR or radio flux.  Using \\textit{Spitzer} FIR data and radio data from the {Very Large Array} ({VLA}) and {Giant Metre-Wave Radio Telescope} ({GMRT}), we quantify the FIR--Radio Correlation as a function of redshift in massive galaxies selected from a near-infrared (NIR) survey of the Extended Chandra Deep Field South (ECDFS).  We use Equation~\\ref{eqn:q_ir} to define the `$q$' index, which quantifies the FRC as the logarithmic ratio between a monochromatic FIR flux ($S_{\\nu,\\text{IR}}$, e.g. at 24, 70 or 160\\mum), and 1.4-GHz radio flux ($S_{\\nu,\\text{1.4~GHz}}$). \\begin{equation} \\label{eqn:q_ir} q_\\text{IR} = \\log_{10}\\left(\\frac{S_{\\nu,\\text{IR}}}{S_{\\nu,\\text{1.4~GHz}}}\\right) \\end{equation} We also investigate the effects of using different FIR bands to quantify the FRC, and the effects of assumptions about the SEDs of the galaxies in the sample.  We employ a `stacking' methodology to recover sufficient signal-to-noise ratios on faint objects to obtain measurements of the average properties of the sample.  The data are described in Section~\\ref{sec:data}, while the binning and stacking methodologies are described in Section~\\ref{sec:method}.  The analysis of SEDs and application of $K$--corrections is covered in Section~\\ref{sec:kcorrs}, and the results are analysed and discussed in Section~\\ref{sec:discussion}.  A concordance cosmology of $\\Omega_m=0.27$, $\\Omega_\\Lambda=0.73$, $H_0=71$\\ kms$^{-1}$Mpc$^{-1}$ is assumed throughout. ",
        "conclusions": "We have studied the FRC as a function of redshift for NIR-selected massive galaxies in the ECDFS, a sample which is unbiased by star-formation activity.  We used a stacking analysis to evaluate the ratios of median FIR/radio fluxes of all galaxies in the sample, divided into redshift bins. This technique traces the typical objects in the population of massive galaxies from low redshift back to their formation epoch.  A thorough analysis of clustering of the sample was used to correct for the differential effects of confusion in the three FIR bands.  $K$--corrections were derived in the radio and FIR using ratios of observed fluxes, ensuring as much as possible a self-consistent analysis.  A mass-limited sub-sample was also stacked to confirm the robustness of the results to Malmquist bias.  The results for $q_{24}$, $q_{70}$ and $q_{160}$ show a slight decline in the \\textit{observed} relations, not dissimilar to the results of previous studies, which can be largely accounted for by the FIR $K$--correction using an M51 template.  After $K$--correction $q_{70}$ is the only monochromatic index to show signs of evolution, suggesting that the 70-\\mum\\ $K$--correction may be less effective as a result of a steep slope in the SED from $\\sim 25-35\\mum$ (corresponding to $z\\sim1-2$) compared with M51.  Observed MIPS colours at all redshifts are more consistent with the M51 template compared with hotter starburst galaxy templates, indicating that the typical IR SEDs of stellar-mass-selected galaxies at redshifts up to $\\sim 0.8$ (at least) appear to be dominated by cold dust.  At higher redshifts it is not possible to constrain the dust temperature with MIPS colours, although it is still clear that M51 is the closest template.  In contrast to this, both radio and total IR luminosities rise significantly with increasing redshift, as do derived SFRs.  Specific SFRs similarly rise steeply, in agreement with {results in the literature \\citep{cowie_new_1996,madau_star_1998,brinchmann_mass_2000, bauer_specific_2005,feulner_connection_2005,prez-gonzlez_stellar_2008, dunne_star_2009,damen_evolution_2009,pannella_star_2009,oliver_specific_2010}.}  The stacked radio data reveal tentative evidence for an evolution in radio spectral index across the redshift range, an unexpected result that implies some change in the radio loss processes in our sample towards higher redshifts.  The most likely explanation seems to be a shift towards greater inverse-Compton losses of the CR electrons at $z>1$, {supporting} the predictions of \\citet{lacki_physics_2009}.  Overall our results show evidence that the FRC, measured from 24-\\mum\\ fluxes or 160-\\mum\\ fluxes closer to the FIR peak, remains roughly constant up to $z\\sim2$, corresponding to 10 Gyr of cosmic time.  This is similar to the conclusions of recent studies including \\citet{ibar_exploringinfrared/radio_2008, garn_relationship_2009, younger_millimetre_2009, ivison_blast:far-infrared/radio_2009} and \\citet{sargent_vla-cosmos_2010}.  The issue is clouded however by measurements at 70\\mum, which appear to show a declining $q$ index with redshift, and when combined into a total IR luminosity, likewise show a slight decline (at low significance).  This most likely implies a steeper spectral slope at wavelengths around $25-35$\\mum\\ (compared with the M51 template), leading to insufficient 70-\\mum\\ $K$--corrections.  But a true evolution in the ratios of 70-\\mum/radio luminosity and of TIR/radio luminosity is plausible, considering the apparent increase in electron-calorimetry behaviour at $z>1$, and considering the fact that rest-frame 24, 70 and 160-\\mum\\ fluxes can arise from different components of the dust in a galaxy.  It is also consistent with the results of \\citet{seymour_investigatingfar-ir/radio_2009} for $q_{70}$ and \\citeauthor{ivison_blast:far-infrared/radio_2009} (2010a/b) using BLAST/Herschel and \\textit{Spitzer} observations to measure $q_\\text{TIR}$.  Constraining the FIR SED is one of the greatest problems in understanding the FRC and the dust emission in general from high redshift star-forming galaxies.  Upcoming surveys with \\textit{Herschel}, such as the \\textit{Herschel} Multi-tiered Extragalactic Survey \\citep[HerMES;][]{oliver_herschel_2010} and the \\textit{Herschel} Astrophysical Terahertz Large Area Survey \\citep[H-ATLAS;][]{eales_herschel_2010} are anticipated to revolutionise our understanding of these topics by providing deep and wide observations spanning the peak of FIR emission across the history of cosmic star formation."
    },
    "1005/1005.4891_arXiv.txt": {
        "abstract": "We present  the  gamma-ray data of the extraordinary flaring activity above 100 MeV from the flat spectrum radio quasar 3C~454.3 detected by AGILE during the month of December 2009. \\cp, that has been among the most active blazars of the FSRQ type since 2007, was detected in the gamma-ray range with a progressively rising flux since November 10, 2009. The gamma-ray flux reached a value comparable with that of the Vela pulsar on December 2, 2009. Remarkably, between {\\es December 2 and 3, 2009} the source more than doubled its gamma-ray emission and became the brightest gamma-ray source in the sky with a {\\es peak flux} of $F_{\\gamma,p} = (2000 \\pm 400) \\times 10^{-8} \\rm ph \\, cm^{-2} \\, s^{-1}$ for a 1-day integration above 100 MeV. The gamma-ray intensity decreased in the following days with the source flux remaining at large values near $F_{\\gamma} \\simeq (1000 \\pm 200) \\times 10^{-8} \\rm ph \\, cm^{-2} \\, s^{-1}$ for more than a week. This exceptional gamma-ray flare dissipated among the largest ever detected intrinsic radiated power in gamma-rays above 100 MeV {\\esc ($L_{\\gamma, source, peak} \\simeq 3 \\times 10^{46} \\; \\rm erg \\, s^{-1} $}, for a relativistic Doppler factor of $\\delta \\simeq 30$). The total isotropic irradiated energy of the month-long episode {\\es in the range 100 MeV -- 3 GeV} is $E_{\\gamma,iso} \\simeq 10^{56} \\, \\rm erg$. We report the intensity and spectral evolution of the gamma-ray emission across the flaring episode. We briefly discuss the important theoretical implications of our detection. ",
        "introduction": "Blazars (a special class of Active Galactic Nuclei with the relativistic jet pointing towards the Earth) show variability across their emitted spectrum on timescales of days, months, years. Rarely, intense gamma-ray flares are detected from blazars with fluxes reaching values near that of the Vela pulsar (i.e., the brightest steady gamma-ray source in the sky with a flux of $F_{\\gamma,Vela} \\simeq 900 \\times 10^{-8} \\rm ph \\, cm^{-2} \\, s^{-1}$ above 100 MeV). Even more rarely, a blazar gamma-ray ``super-flare'' reaches intensities substantially larger than $F_{\\gamma,Vela}$, {\\es as in the case of the June, 1995 flare from PKS 1622-29 (\\cite{1997ApJ...476..692M}). On these rare occasions, the gamma-ray sky is remarkably dominated by a single transient source.  In this \\textit{Paper}, we report the observations by the AGILE satellite of the most recent gamma-ray super-flare from the blazar \\c during the period mid-November/mid-December, 2009. During a 1-month period, this source repeatedly reached a flux near $F_{\\gamma,Vela}$ for about 2 weeks, and then produced a very intense super-flare on {\\es December 2-3}, 2009, with $F_{\\gamma} > 2 F_{\\gamma,Vela} $. This flare turns out to be even more intense than that detected from PKS 1622-29 (\\cite{1997ApJ...476..692M}), and then qualifies as the most intense gamma-ray flare ever observed from a cosmic source at energies above 100 MeV.   \\begin{figure*} % \\begin{center} \\includegraphics[width=12.cm, angle=270]{3c454_flare_agile_paper.eps} \\caption{Gamma-ray emission above 100 MeV from 3C~454.3 as monitored by AGILE. \\textit{(Upper panel:)} the 2 and 1/2 year flux lightcurve covering the period July, 2007 - December, 2009. The black data points are obtained with AGILE in the pointing mode (see \\cite{Vercellone2010:3C454}); the blue data points were collected in spinning mode. \\textit{(Lower panel:)} gamma-ray lightcurve obtained for the period November 07, 2009 and January 9, 2010 (spinning mode). All flux values are obtained with the AGILE standard maximum likelihood analysis, using the {\\es FM3.119} calibrated filter, with standard event selection that takes into account the SAA passage and the Earth albedo photons. The temporal bin for the blue data points is 2 days, except for the 9-days interval centered around the peak, for which a 1-day integration was performed. {\\escc Data obtained with a maximum off-axis angle $\\theta_{m}=60$ degrees.}} \\label{LC} \\end{center} \\end{figure*}    The blazar \\c (PKS~2251$+$158; $z=0.859$) has been extensively studied over the last two decades. A wealth of multifrequency observations were obtained especially after the EGRET detections above 100 MeV during the 90's in the range $F_{\\gamma} = (40-140) \\times 10^{-8} \\rm ph \\, cm^{-2} \\, s^{-1}$ (\\cite{Hartman1992:3C454iauc}; \\cite{Aller1997:3C454_EGRET}). The source entered an active phase in 2000, and was very active in 2005-2006. {\\esc During the May-June, 2005 period the source showed} the strongest optical flare ever recorded in May, 2005, {\\esc and reached the}  optical magnitude  $R \\simeq 12$ (\\cite{vil06, Fuhrmann2006:3c454:opt}).  X-ray (\\citealt{Giommi2006:3C454_Swift}) and hard X-ray observations (\\cite{Pian2006:3C454_Integral}) {\\esc covering this active phase in 2005} detected large fluxes between 10 and 100 mCrab. \\c has been subsequently monitored very extensively at all wavelengths \\citep{vil06,vil07, 2009A&A...504L...9V,rai07,rai08a}. Starting with the July, 2007 AGILE detection above 100 MeV (\\cite{Vercellone2008:3C454_ApJ}), \\c has been very active in gamma-rays, and certainly it can be referred as the most active blazar during the last 2 and 1/2 years. A series of papers describe the AGILE gamma-ray observations in 2007-2009 showing repeated flares usually in coincidence with periods of intense optical and X-ray enhanced activity (\\cite{Vercellone2008:3C454_ApJ}, \\cite{Vercellone2008:3c454:ApJ_P1}, \\cite{Vercellone2010:3C454}, \\cite{2009ApJ...707.1115D}); {\\esc see also the 1AGL catalog (\\cite{2009A&A...506.1563P}}). The multi-year optical evolution of \\c has been presented in \\cite{2008A&A...491..755R}. Also \\textit{Fermi} detected several gamma-ray flaring episodes above 100 MeV (\\cite{2009ATel.2200....1H,2008ATel.1628....1T}), and determined an average spectrum for the August-September, 2009 period in the range 200 MeV -- 10 GeV showing a distinct {\\es break in the power-law spectrum at energies of 2 -- 3 GeV} (\\cite{Abdo2009:3C454}, {\\es \\cite{2010ApJ...710.1271A}}). ",
        "conclusions": "\\c reveals itself as the most prolific gamma-ray blazar during the last 3 years, and is dominating the {\\esc gamma-ray} sky above 100 MeV since mid-2007. During the period of December, 2009 the source showed a dramatic activity reaching and maintaining for several weeks a flux above 100 MeV comparable or larger than the brightest persistent gamma-ray source such as the Vela pulsar. During the period Dec. 2-3, 2009 \\c produced a super-flare that turns out to be the brightest blazar emission episode above 100 MeV ever detected. The AGILE satellite followed {\\esc in a continuous way} the daily evolution of the flaring activity of \\cp. Even though a comprehensive picture of the physical mechanism at work can be obtained only from a multifrequency collection of simultaneous data, restricting ourselves to the gamma-ray range in any case provides very important information on the physics of the source. For a detailed theoretical modelling and a broad spectral evolution of the \\c exceptional activity see \\cite{2010arXiv1005.3263P}."
    },
    "1005/1005.3649_arXiv.txt": {
        "abstract": "We present radio images of a sample of six Wide-Angle Tail (WAT) radio sources identified in the ATLAS 1.4\\,GHz radio survey, and new spectroscopic redshifts for four of these sources. These WATs are in the redshift range of 0.1469$-$0.3762, and we find evidence of galaxy overdensities in the vicinity of four of the WATs from either spectroscopic or photometric redshifts. We also present follow-up spectroscopic observations of the area surrounding the largest WAT, S1189, which is at a redshift of $\\sim$0.22.  The spectroscopic observations, taken using the AAOmega spectrograph on the AAT, show an overdensity of galaxies at this redshift.  The galaxies are spread over an unusually large area of $\\sim$12 Mpc with a velocity spread of $\\sim$4500 km s$^{-1}$. This large-scale structure includes a highly asymmetric FRI radio galaxy and also appears to host a radio relic. It may represent an unrelaxed system with different sub-structures interacting or merging with one another. We discuss the implications of these observations for future large-scale radio surveys. ",
        "introduction": "\\label{intro}  Wide-Angle Tail (WAT) galaxies are radio galaxies whose radio jets appear to bend in a common direction. They are generally detected in dynamical, non-relaxed clusters of galaxies \\citep[e.g.][]{Burns90} and may be used as probes or tracers for clusters \\citep{Blanton00, Blanton01}. Clusters of galaxies are the largest gravitationally bound structures in the Universe and are powerful testbeds of cosmological models \\citep[e.g.][] {Borgani04, Sahlen09, Kravtsov09}. Clusters also host diffuse radio emission in the form of radio haloes and relics \\citep{Giovannini00, Feretti05, Ferrari08, Giovannini09}.  The bent nature of WATs has commonly been attributed to strong intra-cluster winds caused by dynamical interactions such as cluster-cluster mergers \\citep{Burns98}. WATs are preferentially found in enhanced X-ray regions \\citep{Pinkney00} and are usually associated with dominant cluster galaxies \\citep{Owen76}. \\citet{Mao09a} found the tailed radio galaxies, including WATs, to be located in the densest regions of clusters in the local Universe, consistent with earlier studies \\citep[e.g.][] {Burns90, Blanton00, Blanton01}. Thus WATs represent valuable tracers of high density regions in the intracluster medium (ICM), and this approach has been used in a number of recent studies \\citep[e.g.][] {Blanton00, Blanton03, Smolcic07, Giacintucci09, Kantharia09, Oklopcic10}.  Here we present the radio properties of six WATs that we have identified in ATLAS, the Australia Telescope Large Area Survey, carried out with the Australia Telescope Compact Array (ATCA) at 1.4\\,GHz \\citep{Norris06,Middelberg08}.  ATLAS \\footnote{http://www.atnf.csiro.au/research/deep/index.html} will image seven square degrees of sky over two fields to an rms sensitivity of 10 $\\mu$Jy beam$^{-1}$. The ATLAS fields have been observed with a number of different ATCA configurations, and the typical resolution of the observations is $\\sim$10 arcsec.  The two ATLAS fields, Chandra Deep Field South (CDFS) and European Large Area ISO Survey-South 1 (ELAIS-S1), were chosen to coincide with the \\textit{Spitzer} Wide-Area InfraRed Extragalactic (SWIRE) survey program \\citep{Lonsdale03} so that corresponding optical and infrared photometric data are available.  In addition to the radio properties we present new spectroscopic redshifts for four of the WATs and follow-up spectroscopic observations of galaxies in the vicinity of the largest WAT in order to probe its surrounding structure. This WAT was first identified as radio source S1189 by \\citet{Middelberg08}, and is associated with the SWIRE source SWIRE4\\_J003427.54-430222.5 \\citep{Lonsdale03}.  In this paper we present a summary of the data in Section 2, while the WATs in ATLAS are presented in Section 3. Section 4 presents the results of spectroscopic observations of S1189 and its surrounding region, and discusses the large-scale structure in its vicinity.  In Section 5 we discuss cosmological inverse-Compton quenching and the implications for deep wide radio surveys and ATLAS. This paper uses H$_0$ = 71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega$$_M$ = 0.27 and $\\Omega$$_\\Lambda$ = 0.73 and the web-based calculator of \\citet{Wright06} to estimate the physical parameters. Vega magnitudes are used throughout. ",
        "conclusions": "We have identified a sample of six Wide-Angle Tail (WAT) radio sources. We present new spectroscopic redshifts for four of these sources, and find that these WATs lie in the redshift range 0.1469$-$0.3762. We have examined the fields using both spectroscopic and photometric redshifts of galaxies in the vicinity of the WATs and find evidence of an overdensity of galaxies in four of these WATs.  From a more detailed study of the field around S1189 we find an overdensity of galaxies which is spread over $\\sim$12 Mpc and has a velocity spread of $\\sim$4500 km s$^{-1}$, and a velocity dispersion of $\\sim$870 km s$^{-1}$.  This large-scale structure hosts a putative cD galaxy with, at best, weak radio emission, a radio relic which has a size of $\\sim$274 kpc, and an asymmetric FRI radio galaxy with an extent of $\\sim$559 kpc. The peak brightness at the extremities of the outer lobes of the FRI source differ by a factor of $\\sim$4, possibly due to differences in the environment on opposite sides.  The minor axis of the relic is not directed towards either the host galaxy of the WAT or the putative cD galaxy.  This large-scale structure may represent an unrelaxed system with different sub-structures interacting or merging with one another.  Therefore, deep X-ray observations of the field would be very valuable to further understand this interesting large-scale structure.  WATs are known to occur in clusters of galaxies, and could in principle be useful tracers of clusters at moderate and high redshifts. IC cooling of electrons by interaction with CMBR increases rapidly with z. However, this does not imply a sharp drop in the number of WATs at high z.  Deep and wide-field surveys, such as the Evolutionary Map of the Universe (EMU) \\citep{Norris09}, should provide additional information and insights on the range of structures at moderate and high redshifts. We expect these to be invaluable probes of large-scale structure."
    },
    "1005/1005.3822_arXiv.txt": {
        "abstract": "{{We use deep, five band (100--500 $\\mu$m) data from the {\\it Herschel} Lensing Survey (HLS) to fully constrain the obscured star formation rate, SFR$_{\\rm FIR}$, of galaxies in the Bullet cluster ($z$ = 0.296), and a smaller background system ($z$ = 0.35) in the same field. {\\it Herschel} detects 23 Bullet cluster members with a total SFR$_{\\rm FIR}$ = 144 $\\pm$ 14 M$_{\\sun}$ yr$^{-1}$. On average, the background system contains brighter {far-infrared} (FIR) galaxies, with $\\sim$50\\% higher SFR$_{\\rm FIR}$ (21 galaxies; 207 $\\pm$ 9 M$_{\\sun}$ yr$^{-1}$). SFRs extrapolated from 24 $\\mu$m flux via recent templates (SFR$_{24{\\mu}m}$) agree well with SFR$_{\\rm FIR}$ for $\\sim$60\\% of the cluster galaxies. In the remaining $\\sim$40\\%, SFR$_{24{\\mu}m}$ underestimates SFR$_{\\rm FIR}$ due to a significant excess in observed $S_{100}/S_{24}$ (rest frame $S_{75}/S_{18}$) compared to templates of the same FIR luminosity.}} ",
        "introduction": "In the last decade many studies have attempted to quantify the star formation rate (SFR) within cluster galaxies. Ultraviolet and optical observations have successfully identified trends between unobscured star formation and local environment, suggesting that star formation in cluster core galaxies is generally more quenched \\citep[e.g.][]{kod04-1103,por07-1409}. However, star formation can be obscured by dust, which re-emits stellar light in the far-infrared (FIR), peaking at a rest frame $\\lambda_0$ $\\sim$ 100 $\\mu$m. Mid-infrared surveys \\citep[e.g.][]{met05-425,gea06-661,fad08-9} have explored obscured star formation by estimating total FIR luminosity from  template spectra. These templates are often based on small numbers of well constrained local galaxies, e.g. \\citet{rie09-556}.  The PACS \\citep{pog10} and SPIRE \\citep{gri10} instruments, onboard the ESA {\\it Herschel} Space Observatory \\citep{pil10}, enable unprecedented multi-band coverage of the FIR. The {\\it Herschel} Lensing Survey ({HLS}; PI: E Egami) consists of 5-band observations (100--500 $\\mu$m) of 40 nearby clusters ($z$ $\\sim$ 0.2--0.4). Nominally devised to exploit the gravitational lensing effect of massive clusters to observe high redshift galaxies (see \\citealt{ega10} for details on survey design), a useful by-product is deep FIR observations of the clusters themselves. At these redshifts, {\\it Herschel} photometry spans the peak of the dust component, allowing an accurate constraint of far infrared luminosity, {$L_{\\rm FIR}$}, and hence obscured SFR.  During the {\\it Herschel} Science Demonstration Phase, {HLS} observed the Bullet cluster (1E0657--56; $z$ = 0.296). The reason for this choice was two-fold. First, previous studies report bright {submillimeter} galaxies in the background \\citep[e.g. ][]{rex09-348}, {with HLS analysis presented in \\citet{rex10}}. Second, the Bullet cluster is a recent collision of two clusters \\citep{mar02-27}, offering a unique laboratory for the study of star formation within a dynamic environment. The sub-cluster has conveniently fallen through the main cluster perpendicular to the line of sight ($< 8^\\circ$ from the sky plane; \\citealt{mar04-819}). Analysis of X-ray emission shows that a supersonic bow shock precedes the hot gas, while the weak lensing mass profile indicates that this X-ray bright component lags behind the sub-cluster galaxies due to ram pressure \\citep{mar02-27,bar02-816}. A recent mid-infrared study by \\citet{chu09-963} concluded that ram pressure from the merger event had no significant impact on the star formation rates of nearby galaxies. We can re-evaluate these previous studies by using {\\it Herschel} data to constrain {$L_{\\rm FIR}$} directly. In this letter, we present an exploration of obscured star formation in this cluster environment.  \\begin{figure} \\centering \\includegraphics[height=80mm,angle=270,clip]{figs/zhist.eps} \\caption{Distribution of spectroscopic redshifts (0.27 $<$ $z$ $<$ 0.37) for galaxies within the Bullet cluster field (outline). {\\it Herschel} detected galaxies are also shown (filled). In addition to the Bullet cluster ($z$ = 0.296), there is a background {system} at $z$ = 0.350. Dotted lines show our membership limits of 3000 km s$^{-1}$ and 2000 km s$^{-1}$ respectively.} \\label{fig:zhist} \\end{figure} ",
        "conclusions": "Using deep {\\it Herschel} observations (100--500 $\\mu$m) to fully constrain the FIR component, we derive obscured SFRs for galaxies in the Bullet cluster ($z$ = 0.296), and a background system ($z$ = 0.35) in the same field. {{\\it Herschel} detects 23 Bullet cluster members, with a total SFR$_{\\rm FIR}$ = 144 $\\pm$ 14 M$_{\\sun}$ yr$^{-1}$, while the background system contains 21 detections but $\\sim$50\\% higher SFR (207 $\\pm$ 9 M$_{\\sun}$ yr$^{-1}$). The relative distributions of SFR$_{\\rm FIR}$ and optical flux suggest a difference in dust retention between the two systems. For $\\sim$60\\% of galaxies, SFR$_{\\rm FIR}$ agrees well with estimated SFRs from 24 $\\mu$m flux via recent templates. However, the remaining galaxies display a significant excess at 100 $\\mu$m ($\\lambda_{\\rm 0}$ $\\approx$ 75 $\\mu$m) compared to templates, which causes an under-prediction in SFR$_{24{\\mu}m}$}. We note that such an excess is not found in the high redshift, field sample \\citep{rex10}. Future studies will exploit the full range of 5-band {\\it Herschel} cluster observations available in {HLS}, to form a more complete understanding of the environmental effect on obscured star formation rates, and explore the origin and dependencies of the 100 $\\mu$m excess."
    },
    "1005/1005.3539_arXiv.txt": {
        "abstract": "The origin of the atmosphere of the largest moon of Saturn, Titan, is poorly understood and its chemistry is rather complicated. Ground-based millimeter/sub-millimeter heterodyne spectroscopy resolves line shapes sufficiently to determine information in Titan's atmospheric composition (on vertical profiles and isotopic ratios). We test the capabilities of the Swedish Heterodyne Facility Instrument (SHFI), Receiver APEX-1, together with the Atacama Pathfinder EXperiment APEX 12-m telescope for Titan's atmospheric observations. In particular we present sub-millimeter observations of the CO(2-1) and HCN(3-2) lines of the Titan stratosphere with APEX, and with SHFI taken during the Science Verification (SV) instrument phase on March and June 2008. With the help of appropriate radiative transfer calculations we investigate the possibility to constrain the chemical concentrations and optimize the performance of the APEX-1 instrument for inferring vertical profiles of molecular components of the atmosphere of Titan. \\\\ This study attempts to contribute to constrain radiative transfer and retrieval algorithms for planetary atmospheres, and to give hints to the current and future ground and space-based data acquisition leading to a more thorough understanding of the chemical composition of Titan. ",
        "introduction": "Titan's atmosphere exhibits a complex photochemistry. The origin of carbon monoxide (CO) is not well understood (whether photochemical or primordial). Hydrogen Cyanide (HCN), the most abundant nitrile in Titan, is a key intermediate in production of more complex hydrocarbons and organic molecules. HCN in the Titan atmosphere has been discovered by the infrared observations of Voyager\\,1\\textcolor{black}{\\cite{h81}}, and CO has been detected by the ground-based near-infrared observations\\textcolor{black}{\\cite{l83}}. Following these detections, the molecular concentrations of CO and HCN in Titan's atmosphere have been determined from infrared, millimeter, and sub-millimeter observations as well as from modeling (see Tables\\,2 and 3 in section\\,4 for details). CO seems to be uniformly mixed in Titan's atmosphere up to high altitudes. HCN abundances, however, display a steeper profile with ambiguous enrichment values (Table\\,3). Already several observational data confirm that Titan's atmospheric composition is indeed seasonally\\textcolor{black}{\\cite{roe04,teanby08}}, and spatially dependent\\textcolor{black}{\\cite{teanby09}}. Full behavior of the seasonal characteristics of the spatial distribution is not constrained yet due to the limitation in the temporal coverage of the previous observations, and to the use of different instruments. New sub-millimeter observations are required not only to provide new abundance constraints and shed more light on the rate of these seasonal variations, for example, but for support, complement and cross-calibration the ESA's Herschel Space Observatory mission. One of the goals of the Key Program of Herschel, entitled \\textit{Water and Related Chemistry in the Solar System}\\textcolor{black}{\\cite{ha09}}, is to understand water inventory in the Titan's atmosphere as well as distributions of other hydrocarbons and nitriles. During the guaranteed time, line surveys on Titan at the frequency range of 500 and 5000\\,GHz are going to be carried out using two low-to-medium resolution spectrometers (Photodetector Array Camera and Spectrometer (PACS) and Spectral and Photometric Imaging Receiver (SPIRE)). Our model calculations of the synthetic spectra of Titan expected to be observed with Herschel show that several CO and HCN lines will be detected with PACS and SPIRE\\textcolor{black}{\\cite{re09}}. Because of their low-to-medium spectral resolutions, SPIRE and PACS are not capable of measuring the shapes of CO and HCN lines, which prevents us from determining precisely the vertical profile of CO and HCN mixing ratios. Therefore, high spectral resolution  observations of CO and HCN with ground-based sub-millimeter telescopes would significantly improve the accuracy of retrieving the CO and HCN profiles.  The Atacama Pathfinder EXperiment (APEX) 12-m telescope is operational since 2005\\textcolor{black}{\\cite{g06}}, and has been already used to the planetary science (monitoring mesospheric winds of Venus\\textcolor{black}{\\cite{l08}}), among numerous galactic and extra-galactic objects. In spring 2008, the new APEX Swedish Heterodyne Facility Instrument (SHFI) has been commissioned on the APEX\\,12-m telescope\\textcolor{black}{\\cite{vas08,nys09}}, which consists of four single-pixel receivers (APEX-1, APEX-2, APEX-3, and APEX-T2) mounted in a single cryostat located in the Nasmyth\\,A cabin of APEX. In this report we present a summary of our observations of CO and HCN performed with APEX-1 during the Science Verification (SV) phase with the aims to demonstrate the capabilities of APEX and SHFI for Titan's atmospheric observations, and to investigate the possible retrieval of CO and HCN abundances. We also examine the possible detection of the millimeter/sub-millimeter rotational transition of C$_{2}$H$_{2}$ (already observed in the mid-infrared spectra by the Composite Infrared Spectrometer, CIRS, on Cassini\\textcolor{black}{\\cite{vina07}}), and the search for one new compound, HC$_{5}$N(83-82) which the possible presence has been suggested by experimental results simulating Titan's atmosphere \\textcolor{black}{\\cite{van95,coll95,co06}}.  The observations are described in detail in section\\,2 while the radiative transfer modeling and the data analysis are discussed in sections\\,3 and 4, respectively. A brief discussion is given in section\\,5 and a summary in section\\,6. ",
        "conclusions": "How does one explain the behavior of the retrieved HCN vertical profile obtained here? the relative poor signal-to-noise of the observations obtained here results in a poor fit to the HCN line wings. Furthermore, comparing the results obtained with previous results derived from disk-averaged information, interferometer measurements, and space-based observations summarized in Table\\,3 must be considered cautiously. In other words, direct comparison of the results from different instruments requires us to assume that all instrument-related offsets have been accounted for.  Certainly, Titan's atmospheric temperature and composition are not in a steady state. Does the HNC abundance retrieved here perhaps indicate that the composition of Titan's atmosphere has changed during the last years? Fig.\\,7 in Teanby et al.\\,2009\\textcolor{black}{\\cite{teanby09}} shows that the HCN abundance measured by Cassini (at southern latitudes) has not significantly changed between 2006 and 2008. Then it seems to be that possible detections of time variations (if any) between different observations could be maybe due to observing more of the northern hemisphere than of the southern one as Titan's season changes. As we derived disk-averaged information we need to keep in mind, however, that our vertical distributions are mostly representative of the equatorial region of Titan since the measured flux density spectra are more heavily weighted toward equatorial latitudes. In case that vertical distributions would be representative of the northern latitudes, the profile analysis would become rather complicated. There is evidence of a vortex, which acts to separate enriched from unenriched air, at $25-55\\,^{\\circ}$\\,N encircling the North Pole \\textcolor{black}{\\cite{teanby09,fla,ach08}}.  Further higher signal-to-noise HCN observations at different times are necessary in order to open the possibility to detect seasonal and even spatial variations in HCN, and confirm or not an enrichment layer in Titan's stratosphere."
    },
    "1005/1005.3013_arXiv.txt": {
        "abstract": "We demonstrate a new model which uses an ADD type braneworld scenario to produce a multi-state theory of dark matter. Compactification of the extra dimensions onto a sphere leads to the association of a single complex scalar in the bulk with multiple Kaluza-Klein towers in an effective four-dimensional theory. A mutually interacting multi-state theory of dark matter arises naturally within which the dark matter states are identified with the lightest Kaluza-Klein particles of fixed magnetic quantum number. These  states are protected from decay by a combination of a global $U(1)$ symmetry and the continuous rotational symmetry about the polar axis of the spherical geometry. We briefly discuss the relic abundance calculation and investigate the spin-independent elastic scattering off nucleons of the lightest and next-to-lightest dark matter states. \\\\ \\\\ \\\\ ",
        "introduction": "The standard model (SM) of particle physics, while describing the results of collider experiments with unprecedented precision, lacks a suitable candidate for dark matter.  This is a problem that deserves urgent attention, as  cosmological observations have both measured the energy density of atoms and all other SM particles to be $\\Omega_{\\rm SM} h^2 \\simeq 0.02$, and shown that a new type of \\mbox{\\emph{non-SM}} dark matter contributes $\\Omega_{d} h^2 \\simeq 0.11$ to the density of the Universe \\cite{Dunkley:2008ie}. Corroborating evidence from astrophysics and structure formation suggests that dark matter, whatever it is, appears like non-relativistic cold dark matter (CDM) particles interacting at most weakly with SM particles (see, e.g., Refs.~\\cite{Jungman:1995df,Bergstrom:2000pn,Bertone:2004pz,Hooper:2007qk} or, e.g., Ref.~\\cite{D'Amico:2009df} for a pedagogical review).  Most of the theoretical focus to date has been on minimal dark matter theories containing a single new stable elementary particle. However, while this is a reasonable conservative assumption, this need not be the case.  Furthermore, anomalies such as recent cosmic ray results \\cite{:2008zzr,Adriani:2008zr,Abdo:2009zk} and long-observed annual modulation by the DAMA experiment \\cite{Bernabei:2008yi} have sparked theoretical interest in a non-minimal dark sector.  It is worth investigating the possibility that the dark sector might not be so simple and may, in reality, be better described by a multi-particle set up (see, e.g., Refs.~\\cite{Zurek:2008qg,Profumo:2009tb}).  Extra dimensional models have been studied extensively as a possibility for new physics in both particle physics and cosmology. Interest in these models has been fueled by, amongst other things, their ability to accommodate Higgsless electroweak symmetry breaking \\cite{Csaki:2003zu,Csaki:2003sh} and provide potential solutions to the hierarchy problem \\cite{ArkaniHamed:1998rs,ArkaniHamed:1998nn,Antoniadis:1998ig,Randall:1999ee,Davoudiasl:1999tf,Huber:2000ie}. In particular, the universal extra dimensional (UED) model has been a subject of intense investigation recently because it can easily provide a dark matter candidate \\cite{Servant:2002aq,Servant:2002hb,Bertone:2002ms}. This model, originally proposed in Ref.~\\cite{Appelquist:2000nn}, is so named because all particles are allowed to propagate in the bulk and have universal access to all compact dimensions \\cite{Appelquist:2000nn}. This is in contrast to both \\mbox{Arkani-Hamed--Dimopoulos--Dvali} (ADD) type models where all standard model fields are restricted to a brane \\cite{ArkaniHamed:1998rs} or models in which only some of the standard model fields can access the compact dimensions \\cite{Muck:2001yv,Muck:2003kx}. A key issue when constructing a particle theory of dark matter is its stability on cosmological timescales. In the case of the UED models a vestigial discrete translation invariance,  known as Kaluza-Klein (KK) parity, ensures stability of the lightest KK particle (LKP) to all orders in perturbation theory.  In this paper we discuss a theory containing many scalar dark matter states that originate from a single complex scalar in the factorized spacetime, $M_4 \\times S^2$. This spherical compactification, with associated spherical-harmonic eigenfunctions labeled by quantum numbers $\\ell$ and $m$, naturally organizes the associated four-dimensional KK states, $\\tchi_\\ell^m$, into towers of definite $m$, with  $\\tchi_{|m|}^m$ being the lightest state in each tower.  These lightest KK particles are nominally stable due to the residual rotational symmetry about the polar axis of the extra dimensional geometry and an imposed global $U(1)$ symmetry.\\footnote{Stable on cosmologically-interesting timescales.  While gravitational decay may occur we imagine here an effective theory with stabilized extra-dimensions and very massive shape moduli.}    This set of stable KK particles (SKPs) then comprises the set of dark matter states in the theory.  Since this type of compactification does not allow for a chiral zero-mode fermion in the four-dimensional (4D) effective spectrum \\cite{Camporesi:1995fb, Abrikosov:2001nj}, we restrict ourselves to a non-universal extra dimensional scenario.  For simplicity, we also restrict the entire standard model field content, not just the fermions, to the brane on $S^2$. As the position of the standard model brane on $S^2$ is completely arbitrary we can choose our coordinates such that it resides at the north pole.\\footnote{We imagine here that the ultraviolet physics fixing the brane at the pole renders branon fluctuations very massive.}  While the location of the brane breaks some of the spherical rotational symmetry, rotations about the polar axis are preserved. Although we discuss here the simplest spherical case, we expect similar results will hold for any manifold which preserves a continuous isometry of the extra dimensional ground state metric.   In Section~\\ref{sec:formalism} we describe the formalism of a complex scalar field on the $M_4 \\times S^2$ spacetime while Section~\\ref{sec:instability} discusses the instability of the excited KK particles and the stability of the SKPs within our effective framework. In Section~\\ref{sec:abundance} we briefly describe the relic abundances of the corresponding SKPs and in section~\\ref{sec:scat} we investigate the spin-independent nucleon scattering related to relevant direct-detection experiments. We make some final statements and discuss our conclusions in Section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc} We have demonstrated a novel particle model which uses an ADD type braneworld scenario to produce a multi-state theory of dark matter. Compactification of the extra dimensions onto a sphere leads to the association of a single complex scalar bulk field with multiple KK towers in the 4D effective theory. The lightest KK state in each tower was then naturally stabilized by the combination of an imposed $U(1)$ symmetry and a subset of the continuous spherical symmetry  of the extra dimensions(though much of the formalism we discuss here would remain true even if these states were meta-stable on cosmological timescales). We have also shown that viable regions of the parameter space exist that can readily produce the observed dark-matter relic abundance~\\cite{Winslow&Sigurdson}. This model  remains unconstrained after comparison of the spin-independent $\\tchi_0^0$-proton cross section with recent direct detection data. The most important aspect of this paper is that the multi-state nature of the dark matter is a direct and natural consequence of the continuous compactification geometry."
    },
    "1005/1005.3103_arXiv.txt": {
        "abstract": "We present a strong lens analysis of SDSS~J1004+4112, a unique quasar lens produced by a massive cluster of galaxies at $z=0.68$, using a newly developed software for gravitational lensing. We find that our parametric mass model well reproduces all observations including the positions of quasar images as well as those of multiply imaged galaxies with measured spectroscopic redshifts, time delays between quasar images, and the positions of faint central images. The predicted large total magnification of $\\mu \\sim 70$ suggests that the lens system is indeed a useful site for studying the fine structure of a distant quasar and its host galaxy. The dark halo component is found to be unimodal centered on the brightest cluster galaxy and the {\\it Chandra} X-ray surface brightness profile. In addition, the orientation of the halo component is quite consistent with those of the brightest cluster galaxy and member galaxy distribution, implying that the lensing cluster is a relaxed system. The radial profile of the best-fit mass model is in good agreement with a mass profile inferred from the X-ray observation. While the inner radial slope of the dark halo component is consistent with being $-1$, a clear dependence of the predicted A--D time delay on the slope indicates that an additional time delay measurement will improve constraints on the mass model. ",
        "introduction": "SDSS~J1004+4112 is a unique quasar lens system \\citep{inada03,oguri04,inada05,inada08}. A quasar at $z=1.734$ is multiply imaged into five images, with their maximum image separation of $14\\farcs7$, produced by a massive cluster of galaxies at $z=0.68$. It is one of only two known examples of cluster-scale quasar lenses, the other being the triple lens SDSS~J1029+2623 with the maximum image separation of $22\\farcs5$ \\citep{inada06,oguri08}. In addition to the quasar images, SDSS~J1004+4112 contains spectroscopically confirmed multiply imaged galaxies at $z\\sim 3$ \\citep{sharon05}. Both the quasar images and the lensing cluster have been detected in {\\it Chandra} X-ray observations \\citep{ota06}. Moreover, time delays between some of the quasar images have also been measured from long-term optical monitoring observations \\citep{fohlmeister07,fohlmeister08}.  \\begin{table*} \\caption{Previous mass modeling of SDSS~J1004+4112\\label{table:model}} \\begin{center} \\begin{tabular}{lccc} \\hline\\hline Reference & model\\footnotemark[$*$] & constraints (quasar)\\footnotemark[$\\dagger$] & constraints (galaxy)\\footnotemark[$\\dagger$]\\\\ \\hline \\citet{inada03}   & SIE+pert      & pos+flux (4) & $\\cdots$ \\\\ \\citet{oguri04}   & NFW+SIE+pert  & pos+flux (4) & $\\cdots$ \\\\ \\citet{williams04}& non-parametric& pos (4)      & $\\cdots$ \\\\ \\citet{sharon05}  & SPL+gals      & pos (5)      & pos (5)  \\\\ \\citet{kawano06}  & gNFW+SIE+pert & pos+flux (4) & $\\cdots$ \\\\ \\citet{saha06}    & non-parametric& pos (5)      & pos (8) \\\\ \\citet{fohlmeister07} & NFW+deV+gals+pert & pos+flux (5) & $\\cdots$ \\\\ \\citet{saha07}    & non-parametric& pos (5), $\\Delta t$ (2) & pos (8) \\\\ \\citet{inada08}   & NFW+SIE+gals+pert & pos+flux (5), $\\Delta t$ (2) & $\\cdots$\\\\ \\citet{liesenborgs09}& non-parametric & pos (5), $\\Delta t$ (3) & pos (7) \\\\ \\hline This work & gNFW+Jaffe+gals+pert & pos+flux (5), $\\Delta t$ (3) & pos (27)\\\\ \\hline \\multicolumn{4}{@{}l@{}}{\\hbox to 0pt{\\parbox{160mm}{\\footnotesize \\footnotemark[$*$] Name of models: SIE = singular isothermal ellipsoid, pert = external perturbation (e.g., external shear), NFW = Navarro, Frenk \\& White (NFW) profile, SPL = softened power-law model, gals = perturbations from member galaxies, typically modeled by truncated isothermal profiles, gNFW = generalized NFW profile, dev = de Vaucouleurs profile, Jaffe = pseudo-Jaffe profile. \\par\\noindent \\footnotemark[$\\dagger$] ``pos'' indicates constraints from image positions, ``flux'' from fluxes of images, and $\\Delta t$ from time delays. Numbers in parentheses show the number of images used as constraints. }\\hss}} \\end{tabular} \\end{center} \\end{table*}  Such a wealth of observational data available enable detailed investigations of the central mass distribution of the lensing cluster. Indeed, there have been several attempts to model the mass distribution of SDSS~J1004+4112, using either parametric or non-parametric method. \\citet{oguri04} adopted two-component (halo plus central galaxy) model to successfully reproduce the positions of four quasar images, but even the parities and temporal ordering of the quasar images could not be determined because of model degeneracies. \\citet{kawano06} extended mass modeling along this line, and explored how time delay measurements can distinguish different mass profiles. \\citet{fohlmeister07} pointed out that it is important to include perturbations from member galaxies to reproduce the observed time delay between quasar images A and B. On the other hand, \\citet{williams04} and \\citet{saha07} performed non-parametric mass modeling to show possible substructures in the lensing cluster. From the non-parametric mass modeling, \\citet{saha06} and \\citet{liesenborgs09} concluded that the radial mass profile is consistent with that predicted in $N$-body simulation \\citep{navarro97}. We summarize previous mass modeling in Table~\\ref{table:model}.  In this paper, we revisit strong lens modeling of SDSS~J1004+4112 adopting a parametric mass model. We include many observational constrains currently available, including central images of lensed quasars and galaxies, flux rations, and time delay between quasar images (see Table~\\ref{table:model}). In particular this paper represents first parametric mass modeling that includes {\\it both} the quasar time delays and the positions of multiply imaged galaxies as constraints. We compare our best-fit mass model with the {\\it Chandra} X-ray observation of this system \\citep{ota06}.  This paper is organized as follow. We describe our mass model in \\S\\ref{sec:model}. We show our results in \\S\\ref{sec:result}, and give conclusion in \\S\\ref{sec:model}. A new software for gravitational lensing, which is used for the mass modeling, is presented in Appendix. Throughout the paper we adopt the matter density of $\\Omega_M=0.26$ and the cosmological constant of $\\Omega_\\Lambda=0.74$, but regard the dimensionless Hubble constant $h$ as a parameter. With this choice of cosmological parameters, a physical transverse distance of $1h^{-1}$~kpc at the redshift of the lensing cluster ($z=0.68$) corresponds to $0.20$~arcsec. We denote a angular diameter distance from observer to lens as $D_{\\rm l}$, from observer to source as $D_{\\rm s}$, and from lens to source as $D_{\\rm ls}$. ",
        "conclusions": "\\label{sec:summary}  We have revisited parametric mass modeling of SDSS~J1004+4112, a unique quasar--cluster lens system, using a newly developed mass modeling software. We include several new constraints, including positions of spectroscopically confirmed multiply imaged galaxies, time delays between some quasar images, and faint central images. Our model comprising of a halo component modeled by the generalized NFW profile, member galaxies including the brightest cluster galaxy G1, and perturbation terms, well successfully reproduced all observations including time delays.  Unlike earlier claims based on parametric mass modeling, we have found that the center and orientation of the dark halo component is in good agreement with those of member galaxies and {\\it Chandra} X-ray observation, implying that the cluster is highly relaxed. The radial profile from strong lensing is also in excellent agreement with the mass profile inferred from the X-ray observation. Our mass modeling prefers the dark halo component with the inner slope close to $\\alpha=1$, being consistent with so-called NFW density profile. Additional measurement of the time delay between quasar image A and D will be useful to constrain the mass model further. The predicted total magnification of $\\mu_{\\rm tot}=67$ for the NFW profile is quite large compared with those for typical galaxy-scale lenses, because of the shallower density profiles for clusters. This makes the lens system a good site for studying the fine structure of the quasar through microlensing (\\cite{richards04}; \\cite{green06}; \\cite{lamer06}; \\cite{pooley07}) or for studying host galaxies of distant quasars \\citep{ross09}. We note that our result based on parametric mass modeling is broadly consistent with recent non-parametric mass modeling by e.g., \\citet{liesenborgs09}.  Our result suggests that the core of the lensing cluster at $z=0.68$ is highly evolved. Recently, \\citet{limousin10} showed that the cluster MACSJ1423.8+2404 at $z=0.54$ is similarly highly relaxed based on the comparison of mass, light, and gas distributions. Therefore our result may point to the fact that relaxed clusters are quite common already at $z\\sim 0.6$.  It is clear that the lensing cluster SDSS~J1004+4112 is currently one of the best-studied high-redshift clusters whose inner density profile is very tightly constrained by strong lensing. Additional constraints on this cluster with weak lensing, Sunyaev-Zel'dovich effect, and spectroscopic identifications of member galaxies will be important to advance our understanding of high-redshift clusters.  \\bigskip  I would like to thank the referee, Marceau Limousin, for useful comments and suggestion.  \\appendix"
    },
    "1005/1005.1937_arXiv.txt": {
        "abstract": "{ The intermediate-mass star-forming core UYSO~1 has previously been found to exhibit intriguing features. While deeply embedded and previously only identified by means of its (sub-)millimeter emission, it drives two powerful, dynamically young, molecular outflows. Although the process of star formation has obviously started, the chemical composition is still pristine. We present {\\it Herschel} PACS and SPIRE continuum data of this presumably very young region. The now complete coverage of the spectral energy peak allows us to precisely constrain the elevated temperature of 26\\,--\\,28 K for the main bulge of gas associated with UYSO1, which is located at the interface between the hot H {\\sc ii} region Sh 2-297 and the cold dark nebula LDN 1657A. Furthermore, the data identify cooler compact far-infrared sources of just a few solar masses, hidden in this neighbouring dark cloud.  } ",
        "introduction": "Star formation occurs predominantly in structured environments in the case of intermediate and high-mass star formation. In particular, one often finds different stages of evolution located within close proximity. This poses a challenge for working on the earliest phases of star formation, which can only be revealed with long-wavelength observations. To observe these deeply embedded objects at the spectral peak of their emission, airborne or satellite missions are necessary, which in the past delivered only a very modest spatial resolution ($>$20$''$).  We use ESA's new far-IR and sub-millimeter satellite {\\it Herschel} \\citep{A&ASpecialIssue-Herschel} --- with its unprecedented spatial resolution --- to scrutinise a sample of very young low- and high-mass star-forming cores within our programme {``Earliest Phases of Star Formation''} (EPoS, PI: O.~Krause). One of these targets is a core located in the \\object{Canis Majoris OB1/R1} region \\citep[e.g.,][]{2009A&A...506..711G} of the outer Galaxy at a distance of $\\sim 1$~kpc. Originally detected as a distinct submillimetre source in the vicinity of \\object{IRAS 07029-1215} \\citep{2004ApJ...602..843F}, it was coyly named ``UYSO~1'' (unidentified young stellar object 1) since it was only detected at 450 and 850 $\\mu$m and seemed to have no counterpart at other wavelengths.  In the optical, the entire region is dominated by the strong H{\\sc ii} region \\object{Sh 2-297}, excited by the early B star \\object{HD 53623}. The mid-IR emission is dominated by warm dust and PAH emission from the associated southern photon-dominated region (PDR). All these sources, however, do not coincide with UYSO~1, located more than 1$'$ to the north-west. Further to the west, one finds the optical dark cloud \\object{LDN 1657A}, about which little is known. Members of our group carried out a multiwavelength study of \\object{UYSO 1} \\citep{2009A&A...493..547F}, which revealed more puzzles.  The source drives a dynamically young, but strong CO molecular outflow. The imprints of two crossed jets are visible in shocked H$_2$ emission in the near-IR, and their vertex is very close to the position of UYSO~1.  Millimeter continuum interferometry resolved the source into two peaks separated by 4\\farcs2.  {\\it Spitzer}/MIPS data indicated that at 24 and 70 $\\mu$m, UYSO~1 is still too deeply embedded to be directly detected, a finding we revise in the present paper for 70 $\\mu$m with {\\it Herschel}.  In addition to the core being chemically pristine and showing no signs of  more evolved chemistry, this object is very interesting for the investigation of the early evolution of (intermediate-mass) protostars. ",
        "conclusions": "We have presented new {\\it Herschel} PACS and SPIRE scan map data for the star-forming complex containing the intermediate-mass core UYSO~1. They show the dust emission structures associated with this region in unprecedented detail. The high spatial resolution of the PACS 70 $\\mu$m map facilitates the differentiation of the true emission peak position from emission arising from the surrounding PDR material. The PACS data show that the 70 $\\mu$m emission is closely associated with the known location of the central millimetre peak(s) in UYSO~1. This revises an earlier finding of large offsets (based on {\\it Spitzer}/MIPS 70 $\\mu$m data).  Hence, the measured {\\it Herschel} photometry can be used with more confidence than for previous data to estimate the SED of UYSO~1. Since the core was not clearly detected at 24 $\\mu$m (which otherwise would indicate a distinct component of warmer dust), a one-temperature modified black-body component with T = 28 K fits the data between 70\\,--\\,500 $\\mu$m very well. While clearly not in the hot-core/hot-corino regime, this temperature is higher than the typical values of $<$17--20 K found in other very young regions of star formation. The known outflow activity can be one explanation. However, the influence of the nearby PDR has also to be taken into account \\citep{2009A&A...493..547F}. The Herschel data reveal the serendipitous discovery of five compact very red sources, situated to the north-west of UYSO~1, within the neighbouring dark cloud LDN 1657A. These objects have lower temperatures than UYSO~1. A particularly intriguing object is LDN 1657A-4, which has a very red SED and exhibits a 24 $\\mu$m shadow. We note that two different methods for the column density estimation provided very consistent results. For both objects, the \\citet{1994A&A...291..943O} grain opacities seem more suitable than the use of ISM grains."
    },
    "1005/1005.4474.txt": {
        "abstract": "%Studies of cosmic-ray accelerations and interactions in supernova remnants (SNRs) are critical step to solve the mystery of the cosmic-ray origins. %Observation of enhanced $\\pi^0$ gamma rays in the GeV--TeV band from SNRs interacting with molecular clouds is one of the promising approaches to confirm hadronic nature of cosmic rays accelerated in SNRs. We present detailed analysis of the two gamma-ray sources, \\FGLnorth\\ and \\FGLsouth, that have been found toward the supernova remnant~(SNR) W28 with the Large Area Telescope~(LAT) on board the \\emph{Fermi} Gamma-ray Space Telescope. \\FGLnorth\\ is found to be an extended source within the boundary of SNR W28, and to extensively overlap with the TeV gamma-ray source HESS~J1801$-$233, which is associated with a dense molecular cloud interacting with the supernova remnant. The gamma-ray spectrum measured with LAT from 0.2--100 GeV can be described by a broken power-law function with a break of $\\sim$~1~GeV, and photon indices of 2.09~$\\pm$~0.08~(stat)~$\\pm$~0.28~(sys) below the break and 2.74~$\\pm$~0.06~(stat)~$\\pm$~0.09~(sys) above the break. Given the clear association between HESS~J1801$-$233 and the shocked molecular cloud and a smoothly connected spectrum in the GeV--TeV band, we consider the origin of the gamma-ray emission in both GeV and TeV ranges to be the interaction between particles accelerated in the SNR and the molecular cloud. The decay of neutral pions produced in interactions between accelerated hadrons and dense molecular gas provide a reasonable explanation for the broadband gamma-ray spectrum. % \\FGLsouth\\, located outside the southern boundary of SNR W28, is consistent with a point source. \\FGLsouth\\, located outside the southern boundary of SNR W28, cannot be resolved. An upper limit on the size of the gamma-ray emission was estimated to be $\\sim$~16$'$ using events above $\\sim$~2~GeV under the assumption of a circular shape with uniform surface brightness. It appears to coincide with the TeV source HESS~J1800$-$240B, which is considered to be associated with a dense molecular cloud that contains the ultra compact \\hii\\ region W28A2 (G5.89$-$0.39). We found no significant gamma-ray emission in the LAT energy band at the positions of TeV sources HESS~J1800$-$230A and HESS~J1800$-$230C. %The LAT data for HESS~J1800$-$230A combined with the TeV data points %indicate a spectral break around 100~GeV. The LAT data for HESS~J1800$-$230A combined with the TeV data points indicate a spectral break between 10~GeV and 100~GeV.   %Again, the gamma rays are naturally explained by $\\pi^0$ decays. %The W28A2 is a possible energy source although it requires extremely dense gas due to the energetics of cosmic rays. %Cosmic rays escaped from W28 are another plausible explanation given the energetics. %The SNR W28 is a prime target for observations of GeV gamma rays due to the detection of TeV gamma rays and its interaction with molecular clouds along its northeastern boundary and other clouds situated nearby. %We report observations of two GeV gamma-ray sources around W28 with the \\emph{Fermi} Gamma-ray Space Telescope~(FGST) - Large Area Telescope~(LAT) in the energy band between 200~MeV and 100~GeV. %We accumulated high energy gamma-ray events (200~MeV $<$ E $<$ 100~GeV) for $\\sim$~11 months. %Spatial extension and source position, and spectra were investigated by carrying out maximum likelihood fits to the data set.  %One is spatially extended, and is positionally coincident with TeV gamma-ray source HESS~J1801$-$233 and also the molecular clouds interacting with this SNR at the northeast boundary of the SNR. %Decays of $\\pi^0$s produced by the interaction of hadrons accelerated in the SNR naturally explain the origin of gamma-ray emission in this source since electron models faces difficulties with high electron/proton ratio and energetics of cosmic rays constrained by the gamma-ray flux. %W28 is the most plausible energy source due to observational evidence of interaction with the molecular clouds, such as OH masers.  %We found no significant gamma-ray emission in the LAT energy band at the positions of TeV sources, HESS~J1800$-$230A and HESS~J1800$-$230C, and obtained upper limits. %The upper limits for HESS~J1800$-$230A combined with the TeV data points indicate a spectral break around 100~GeV. %The energy-dependent escape of cosmic rays from the SNR and/or propagation effect of cosmic rays may produce such a break.  %HESS J1801-233 and HESS J1800-240, and several % cloud components in projection may contribute to the VHE emission. %The clouds have components covering a broad velocity range encompassing % the distance estimates for W 28 (~2 kpc) and extending up to ~4 % and the \\hii\\ regions M 8 and M 20, along with their associated open % clusters. Further sub-mm observations would be recommended to probe in % detail the dynamics of the molecular clouds at velocites > 10 km s-1 % and their possible connection to W 28. ",
        "introduction": "Diffusive shock acceleration operating at supernova shock waves can distribute particles to very high energies with a power-law form having number index about 2 \\citep[e.g.,][]{blandford87}. It is generally expected that if a dense molecular cloud is overtaken by a supernova blast wave, the shocked molecular cloud can be illuminated by relativistic particles accelerated at supernova shocks \\citep{Aharonian94}. If the accelerated particles are comprised mostly of protons, say $>100$ times more abundant than electrons like the observed Galactic cosmic rays, decays of neutral pions produced in inelastic collisions of the accelerated protons with dense gas are expected to be a dominant radiation component in the gamma-ray spectrum of the cosmic-ray-illuminated molecular cloud. Although an earlier attempt to detect TeV gamma-ray emission from supernova remnants (SNRs) that have evidence for molecular cloud interactions with the Whipple telescope failed \\citep{Whipple98}, two archetypical SNRs interacting with molecular clouds, IC~443 \\citep{Albert07, Acciari09} and W28 \\citep{Aharonian08}, have been detected with the current generation of imaging atmospheric Cherenkov telescopes. However, the identities of the particles responsible for the TeV sources remains elusive.  The advent of the Large Area Telescope (LAT) onboard the \\emph{Fermi} Gamma-ray Space Telescope has brought a new opportunity to study the gamma-ray emission from SNRs at GeV energies. LAT observations of middle-aged SNRs interacting with molecular clouds, W51C~\\citep{LAT-W51C}, W44~\\citep{LAT-W44}, and IC~443~\\citep{LAT-IC443}, have revealed bright extended gamma-ray sources coincident with the SNRs. The gamma-ray luminosity reaches $\\sim 10^{36}\\ \\rm erg\\ s^{-1}$, which immediately rules out an inverse-Compton origin of the GeV gamma rays since it requires total electron energy comparable to or larger than the typical kinetic energy released by a supernova explosion, $\\sim 10^{51}$~erg. The gamma-ray spectra of the three remnants exhibit remarkable spectral breaks at an energy of several GeV, making these SNRs much less luminous at TeV energies. This characteristic demonstrates the importance of observations at GeV energies. %In this paper, detailed analysis of LAT data in a field that contains %an archetypical SNR-cloud system, W28, will be presented.   W28 is a mixed-morphology SNR, characterized by center-filled thermal X-ray emission and shell-like radio morphology. In addition X-ray observations show limb-brightened shells in the northeast and southwest \\citep{Rho02}. The shell-like radio emission is prominent in the northeastern region with slightly fainter emission at the northern boundaries \\citep{Dubner00}. Interactions of the SNR with molecular clouds~\\citep{Wootten81} along its northern and northeastern boundaries are traced by the high concentration of 1720~MHz OH masers \\citep{Frail94, Claussen97,Claussen99}, and  high density ($\\nh > 10^3$~cm$^{-3}$) shocked gas \\citep{Arikawa99, Reach05}. The overall shape of W28 is elliptical with a size of $50'\\times 45'$. W28 is located within a complex, star-forming region along the Galactic plane toward the large \\hii\\ regions (M8 and M20) and young clusters~(e.g., NGC~6530)~\\citep{goudis76}. The observations of molecular lines place SNR W28 at a distance of $\\sim 2$ kpc~\\citep{Velazquez02}. Estimates for its age vary between $3.5$ and $15\\times 10^4$~yrs \\citep{Kaspi93}. W28 is considered to be an evolved remnant in the radiative stage of SNR evolution \\citep{Lozinskaya92}, which is supported by optical observations \\citep{Lozinskaya74}. Measurements with Energetic Gamma-Ray Experiment Telescope (EGRET) onboard the \\emph{Compton Gamma-ray Observatory} found a gamma-ray source, 3EG~J1800-2338~\\citep{Hartman99} in the W28 field. However, its association to SNR W28 remained unclear mainly due to large source location uncertainties from EGRET. A gamma-ray source is listed in the W28 field in the AGILE (Astro-rivelatore Gamma a Immagini LEggero) one year catalog~\\citep{AGILEcatalog}. However, detailed analysis of this field is not published by AGILE yet.      H.E.S.S. observations of the W28 field have revealed four TeV gamma-ray sources positionally coincident with molecular clouds \\citep{Aharonian08}: HESS~J1801$-$233, located along the northeastern boundary of W28, and a complex of sources, HESS~J1800$-$240A, B and C,  located $\\sim 30\\arcmin$ south of SNR W28. HESS~J1801$-$233 coincides with a molecular cloud interacting with SNR W28, providing one of the best examples of a cosmic-ray-illuminated cloud. Understanding the origins of TeV emission in HESS~J1800$-$240ABC is  of particular interest; they may be due to cosmic rays that have diffused from W28.  In this paper, we report \\emph{Fermi} LAT observations of gamma-ray sources in the W28 field in the GeV domain. First, we give a brief description of the observation and gamma-ray selection in Section~\\ref{sec:obs}. The analysis procedures and results are explained in Section~\\ref{sec:ana}, where the spatial extension and spectra of the LAT sources in the W28 field are described. Discussion is given in Section~\\ref{sec:discuss}, followed by conclusions in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have investigated two LAT sources in the W28 field. \\SourceN\\ (\\FGLnorth) which is located at the northeast boundary of the SNR W28 is positionally coincident with shocked molecular clouds, and is spatially extended. The spectrum has a break around 1.0~GeV and smoothly connects to the TeV spectrum, suggesting a physical relationship. % and a spectral curvature due to decays of $\\pi^0$s. } Decay of $\\pi^0$s produced by the interaction of an SNR with molecular clouds naturally explains the gamma rays from \\SourceN\\ based on the spatial correlation between GeV gamma rays and molecular clouds and the energetics of cosmic rays. Electron bremsstrahlung can not be ruled out completely although it requires a low density and low magnetic field in contradiction with the association with the molecular clouds. W28 is the most plausible energy source due to the observational evidence of interaction with the molecular clouds. The soft spectrum of the gamma rays may be explained by the time evolution of non-thermal emission from molecular clouds illuminated by cosmic rays from a nearby SNR due to energy-dependent diffusion of cosmic rays.  %The \\SourceS\\ (\\FGLsouth) is consistent with the emission from a point source and spatially coincides with the TeV source HESS~J1800$-$240B, molecular clouds, and the ultracompact \\hii\\ region W28A2~(G5.89$-$0.39).  The \\SourceS\\ (\\FGLsouth) was found to have no significant extension and spatially coincides with the TeV source HESS~J1800$-$240B, molecular clouds, and the ultracompact \\hii\\ region W28A2~(G5.89$-$0.39). %Again, leptonic models require high electron/proton ratio and the $\\pi^0$ decay is natural explanation of the origin of the gamma rays. The compact \\hii\\ region W28A2 and the SNR W28 are possible energy sources, but the W28A2 hypothesis requires extremely dense gas.  %While no significant LAT counterpart is found at the positions of %HESS~J1800$-$230A and HESS~J1800$-$230C, the LAT upper limits for %HESS~J1800$-$230A coupled with the H.E.S.S. data points imply a %spectral break around 100~GeV. While no significant LAT counterpart is found at the positions of HESS~J1800$-$230A and HESS~J1800$-$230C, the LAT upper limits for HESS~J1800$-$230A coupled with the H.E.S.S. data points imply a spectral break between 10~GeV and 100~GeV. %The energy-dependent escape from SNR and/or propagation of cosmic rays may be responsible for this break.   %our observations with the HESS \u00a1\u00f1?\u00a1\u00f1?-ray telescopes %have revealed VHE \u00a1\u00f1?\u00a1\u00f1?-ray sources in the field of %W 28 which positionally coincide well with molecular clouds. %HESS J1801?233 is seen toward the northeast boundary of %W28, while HESS J1800?240 situated just beyond the southern %boundary ofW28 comprises three components.Our studies with %NANTEN 12CO(J = 1?0) data show molecular clouds spanning %a broad range in local standard of rest velocity VLSR = 5 %to ?20 km s?1, encompassing the distance estimates for W 28 %%and various star formation sites in the region. If connected, and %at a distance ?2 kpc, the clouds may be part of a larger parent %cloud possibly disrupted by W 28 and/or additional objects %related to the active star formation in the region. Cloud components %up to ?4 kpc distance (VLSR > 10 km s?1) however, remain %a possibility. %The VHE/molecular cloud association could indicate a %hadronic origin for the VHE sources in the W 28 field. Under %assumptions of connected cloud components at a common distance %of 2 kpc, or, alternatively, separate cloud components %at 2 and 4 kpc, a hadronic origin for the VHE emission implies %cosmic-ray densities ?10 to ?30 times the local value. %W 28 could provide such densities in the case of slow diffusion. %Additional and/or alternative particle accelerators such %as \\hii\\ regions representing very young stars, other SNRs/SNR %candidates and/or several open clusters in the region may also %be contributors. Alternatively, if cloud components at VLSR > %10 km s?1 are at distances d ? 4 kpc, as-yet undetected particle %accelerators in the Scutum-Crux arm may be responsible. %Detailed modeling (beyond the scope of this paper), and further %multiwavelength observations of this region are highly recommended %to assess further the relationship between the molecular %gas and potential particle accelerators in this complex region, as %well as the nature of the acclerated particles. In particular, further %sub-mm observations (e.g. at high CO transitions) will provide %more accurate cloud mass estimates, and allow to search for %disrupted/shocked gas towards the southern VHE sources. Such %studies will be valuable in determining whether or notW 28 and %other energetic sources have disrupted molecular material at line %velocities >10 km s?1."
    },
    "1005/1005.4981_arXiv.txt": {
        "abstract": "The streaming instability (SI) provides a promising mechanism for planetesimal formation because of its ability to concentrate solids into dense clumps. The degree of clumping strongly depends on the height-integrated solid to gas mass ratio $Z$ in protoplanetary disks (PPDs). In this letter, we show that the magnitude of the radial pressure gradient (RPG) which drives the SI (characterized by $q\\equiv\\eta v_K/c_s$, where $\\eta v_K$ is the reduction of Keplerian velocity due to the RPG and $c_s$ is the sound speed) also strongly affects clumping. We present local two-dimensional hybrid numerical simulations of aerodynamically coupled particles and gas in the midplane of PPDs. Magnetic fields and particle self-gravity are ignored.  We explore three different RPG values appropriate for typical PPDs: $q=0.025, 0.05$ and $0.1$. For each $q$ value, we consider four different particle size distributions ranging from sub-millimeter to meter sizes and run simulations with solid abundance from $Z=0.01$ up to $Z=0.07$. We find that a small RPG strongly promotes particle clumping in that: 1) At fixed particle size distribution, the critical solid abundance $Z_{\\rm crit}$ above which particle clumping occurs monotonically increases with $q$; 2) At fixed $Z$, strong clumping can occur for smaller particles when $q$ is smaller. Therefore, we expect planetesimals to form preferentially in regions of PPDs with a small RPG. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.1089_arXiv.txt": {
        "abstract": "{} {We exploit deep observations of the GOODS-N field taken with PACS, the \\textit{Photodetector Array Camera and Spectrometer}, on board of \\textit{Herschel}, as part of the \\textit{PACS Evolutionary Probe} guaranteed time (PEP), to study the link between star formation and stellar mass in galaxies to $z\\sim 2$.} {Starting from a stellar mass -- selected sample of $\\sim4500$ galaxies with mag$_{4.5 \\mu m}<23.0$ (AB), we identify $\\sim350$ objects with a PACS detection at 100 or 160 $\\mu$m and $\\sim1500$ with only {\\textit Spitzer} 24 $\\mu$m counterpart. Stellar masses and total IR luminosities ($L_{IR}$) are estimated by fitting the Spectral Energy Distributions (SEDs). } {Consistently with other \\textit{Herschel} results, we find that $L_{IR}$ based only on 24 $\\mu$m data is overestimated by a median factor $\\sim1.8$ at $z\\sim2$, whereas it is underestimated (with our approach) up to a factor $\\sim1.6$ at $0.5<z<1.0$. We then exploit this calibration to correct $L_{IR}$ based on the MIPS/\\textit{Spitzer} fluxes. These results clearly show how \\textit{Herschel} is fundamental to constrain $L_{IR}$, and hence the SFR, of high redshift galaxies. Using the galaxies detected with PACS (and/or MIPS), we investigate the existence and evolution of the relations between the star formation rate (SFR), the specific star formation rate (SSFR=SFR/mass) and the stellar mass. Moreover, in order to avoid selection effects, we also repeat this study through a stacking analysis on the PACS images to fully exploit the far-IR information also for the {\\textit {Herschel}} and {\\textit {Spitzer}} undetected subsamples. We find that the SSFR-mass relation steepens with redshift, being almost flat at $z<1.0$ and reaching a slope of  $\\alpha=-0.50^{+0.13}_{-0.16}$ at $z\\sim2$, at odds with recent works based on radio-stacking analysis at the same redshift. The mean SSFR of galaxies increases with redshift, by a factor $\\sim15$ for massive $M>10^{11}M_{\\odot}$ galaxies from $z=0$ to $z=2$, and seems to flatten at $z>1.5$  in this mass range. Moreover, the most massive galaxies have the lowest SSFR at any $z$, implying that they have formed their stars earlier and more rapidly than their low mass counterparts ($downsizing$).  } {} ",
        "introduction": "\\label{intro} The link between galaxy stellar mass and star formation rate (SFR), and its cosmic evolution is crucial to shed light on the processes of galaxy formation. The specific SFR (SSFR = SFR/mass) plays an important role as it measures the star formation efficiency of a galaxy and the fraction of a galaxy mass can be converted into stars per unit time (see e.g. de Cunha et al. 2010). Several studies at $0<z<3$ report similar findings: (1) the SSFR increases with redshift at all masses; (2) the SSFR of massive galaxies is lower at all $z$ (e.g. Feulner et al. 2005, Erb et al. 2006, Perez-Gonzalez et al. 2008, Damen et al. 2009, Dunne et al. 2009). However, the scatter and the exact slopes of these relations are still not clear. In particular, the dependence of SSFR on mass is one of the most debated open questions. On the one hand, radio-stacking analysis of galaxies  at $z=1.5-2.0$ selected in the $K$-band found a weak (or even absent) SSFR-mass correlation, with an indication of steepening at higher redshift (Pannella et al. 2009, Dunne et al. 2009). On the other hand, other results based on UV to mid-IR SFR tracers indicate a clear decrease of the SSFR with increasing mass (e.g. Feulner et al. 2005, Erb et al. 2006, Noeske et al. 2007b, Cowie \\& Barger 2008). Some of these discrepancies could be due on one hand to the effects of dust extinction corrections for the SFR estimates based on UV--optical indicators, and on the other hand to the assumptions and large extrapolations on the infrared SED shape and luminosity adopted for the SFR estimates based on 24$\\mu$m data.  The advent of the {\\it Herschel Space Observatory} (Pilbratt et al. (2010) finally allows us to robustly derive the total infrared (IR) luminosity ($L_{IR}$) of galaxies,  by directly sampling the peak of the thermal emission of dusty galaxies up to $z\\sim3$. In order to place new and stringent constraints on the SSFR evolution, in this work we exploit the PACS Evolutionary Probe (PEP) guaranteed time data collected in the GOODS-North field with the PACS (Poglitsch et al. 2010) instrument. The PEP observations are described in Berta et al. (2010, this Issue Appendix A). Ancillary data, both photometric and spectroscopic,  including UV (GALEX), optical (HST), near-IR (FLAMINGOS, IRAC/{\\textit Spitzer}) and mid-IR (MIPS/{\\textit Spitzer}) data, have been collected to build up a reliable multiwavelength catalog and photometric redshifts. We adopt $h = 0.7$, $\\Omega_{\\Lambda}=0.73$ and $\\Omega_m= 0.27$. ",
        "conclusions": ""
    },
    "1005/1005.3874_arXiv.txt": {
        "abstract": "We present images of extended \\Ha\\ clouds associated with 14 member galaxies in the Coma cluster obtained from deep narrow band imaging observations with the Suprime-Cam at the Subaru Telescope. The parent galaxies of the extended \\Ha\\ clouds are distributed farther than 0.2 Mpc from the peak of X-ray emission of the cluster. Most of the galaxies are bluer than $g-r \\approx 0.5$ and they account for 57\\%\\ of the blue ($g-r<0.5$) bright ($r<17.8$ mag) galaxies in the central region of the Coma cluster. They reside near the red- and blue-shifted edges of the radial velocity distribution of Coma cluster member galaxies. Our findings suggest that the most of the parent galaxies were recently captured by the Coma cluster potential and are now infalling toward the cluster center with their disk gas being stripped off and producing the observed \\Ha\\ clouds. ",
        "introduction": "Clusters of galaxies are ideal laboratories for investigating environmental effects on galaxy evolution. It is well known that galaxies residing in the central region of clusters are generally deficient in gas% \\citep[e.g.,][]{Giovanelli1985, Gavazzi1987, Solanes2001, Boselli2006, BravoAlfaro2009}. This indicates that some gas removal mechanisms work quite efficiently in the cluster center. Recently, \\citet{vanderWel2010} studied the shape and magnitude of quiescent galaxies in SDSS and found that galaxies in dense environments include gas-removed spiral galaxies, adding to the evidence for rapid gas removal. A number of possible physical processes have been proposed to explain efficient gas removal in clusters: ram pressure \\citep{Gunn1972} or viscous stripping \\citep{Nulsen1982}, tidal stripping by galaxy-galaxy interactions \\citep{Toomre1972}, or the cluster potential \\citep{Byrd1990}, evaporation \\citep{Cowie1977}, and to a lesser degree, galaxy harassment \\citep{Moore1996}.  One critical clue for understanding the gas removal is detection of the stripped gas. Recent deep \\Ha\\ imaging observations revealed that some of the galaxies in clusters are accompanied by \\Ha\\  clouds. In the nearby cluster Abell 1367, \\citet{Gavazzi2001} found extended ionized gas regions associated with starburst irregular galaxies. The sizes of the regions are 75$\\times$8 and 50$\\times$8 kpc. \\citet{Yoshida2002} discovered another extended emission line region ($\\sim$35 kpc) in the Virgo cluster, which was found to be a part of 110$\\times$25 kpc of HI gas \\citep{Oosterloo2005}. \\citet{Yagi2007} reported a narrow and long (60$\\times$2 kpc) \\Ha\\ emitting region associated with a galaxy of the Coma cluster. \\citet{Sun2007} found a $\\sim$ 40 kpc \\Ha\\ tail of a galaxy in Abell 3627. The galaxy also had a $\\sim$ 70 kpc X-ray tail. \\citet{Sun2010} investigated the galaxy to show two of $\\sim$ 80 kpc tails, and also reported one more galaxy with $>$20 kpc \\Ha\\ tail in the cluster. \\citet{Kenney2008} detected 120 kpc \\Ha\\ filaments from NGC 4438, a spiral galaxy in the Virgo cluster. \\citet{Yoshida2008} reported a galaxy with detached \\Ha\\ clouds with young blue stars in the Coma cluster, which resembles two galaxies in z$\\sim$0.2 clusters reported by \\citet{Cortese2007}. These \\Ha\\ cloud findings imply that gas removal mechanisms to form such extended emission line regions are common in clusters of galaxies. However, it is not yet clear how common this phenomenon is.  The Coma cluster is the nearest rich cluster of galaxies, and is the best target for investigating the \\Ha\\ clouds in the cluster environment with a deep and spatially resolved observation. We have performed a deep \\Ha\\ narrow band imaging survey of the Coma Cluster using the Suprime-Cam at the Subaru Telescope and have already reported some of the prominent clouds detected \\citep{Yagi2007,Yoshida2008}. Here, we report significant, extended \\Ha\\ clouds around further 12 galaxies, increasing the sample to a total of 14. This is the first paper in a series providing a statistical discussion of extended emission line regions in the central region of the Coma cluster. The nature of the clouds and physical interpretations will be presented in subsequent papers.  We assume that the distance modulus of the Coma cluster is $(m-M)_0=35.05$ \\citep{Kavelaars2000} and ($h_0$,$\\Omega_M$,$\\Omega_\\lambda$)=(0.71,0.27,0.73) \\citep{Larson2010}. Under these assumptions, the angular diameter distance is 97.47 Mpc, and 1 arcsec corresponds to 0.473 kpc. ",
        "conclusions": "\\subsection{Possible evolution sequence of the parent galaxies}  The findings presented in Sections \\ref{sec:result:pos} and \\ref{sec:result:color} suggest that the parent galaxies were recently captured by the Coma cluster potential and are now infalling toward the cluster center, while their disk gas is stripped off and seen as \\Ha\\ clouds.  The classification based on apparent morphology may imply an evolutionary sequence of gas stripping: a star-forming disk galaxy (\\#1) evolves into star-forming + dead disk phase (\\#2) and then the gas clouds are detached (\\#3). Recent numerical simulations of ram-pressure stripping \\citep{Kronberger2008,Kapferer2009,Tonnesen2010} also showed the evolution of \\Ha\\ appearance in $\\sim$ 500 Myr. However, it is puzzling that \\#3 phase galaxies consist of less luminous blue galaxies and bright red giants. Moreover, GMP2374, the brightest \\#1 galaxy, shows a red color. Although part of the red color is due to possible extinction by internal dust, the spectrum of the whole galaxy in the GOLDMINE database \\citep{Gavazzi2003} indicates that it is dominated by intermediate and old stellar populations as inferred from its large D4000 and small CaHK ratio. Therefore it is difficult to assume that this galaxy will evolve into a post-starburst phase with a blue color. The parents may consist of several kinds of galaxies on different evolutionary paths.  As a second parameter of various evolutionary paths, we estimated the masses of the parents. Higher mass galaxies should have larger potential wells and retain some gas for star formation, while lower mass galaxies might simply evolve from type \\#1 to \\#3 without passing through \\#2.  The masses of the parents were estimated by three different methods and are given in Table 4. \\citet{Bell2003} gives stellar mass-to-light ratios (M/L) as a function of color. They adopted a \"diet\" Salpeter initial mass function (IMF), which has 30\\% lower mass than Salpeter IMF \\citep{Salpeter1955}. We adopted \\[ \\log_{10} (M/L)_{\\rm i-band} = 0.006 + 1.114(r-i) - 0.15 \\] and calculated the stellar mass from the SDSS $r$ and $i$ magnitudes. We neglected K-correction for both the $i$ magnitude and $r-i$ colors, which were smaller than 0.1 mag. The possible effects of internal dust were also neglected. We also estimated the masses of the parents from the model SED fitting, using kcorrect v4.1.4 \\citep{Blanton2007} with SDSS and public GALEX magnitudes and redshifts as in \\citet{Miller2009}. The recommended corrections on the kcorrect web page \\footnote{http://howdy.physics.nyu.edu/index.php/Kcorrect} were applied, which included the correction between the AB and SDSS magnitudes, and the minimum uncertainty of the SDSS magnitude. The kcorrect used \\citet{Chabrier2003} IMF and Padova 1994 isochrones \\citep[and references therein]{Bruzual2003}. Yet another mass estimation was available in the MPA-JHU SDSS catalog \\footnote{http://www.mpa-garching.mpg.de/SDSS/DR7/}. The DR7\\_v5.2 version included 11 parents of this study. We adopted the ``MEDIAN'' column as the estimated mass. The IMF by \\citet{Kroupa2001} was adopted for the catalog\\citep{Kauffmann2003}.  Although there were variations among the mass estimations of the three different methods, we can see that the type \\#3 dwarfs had less than $2 \\times 10^9$ M$_{\\sun}$. On the other hand, the classification between \\#1 and \\#2 showed no clear correlation with mass. In addition, the lowest mass \\#3 parents ($<2 \\times 10^9$ M$_{\\sun}$) stopped star formation quickly, while most of the \\#1 and \\#2 parents, the masses of which were $10^9-10^{11}$ M$_{\\sun}$, continued star formation. The correlation of the detachment of the clouds, the activity of star formation, and the mass implied that the quick quench of the star formation in the \\#3 blue dwarfs was due to shallow potentials and/or small amounts of gas. This result regarding the dependency of the star formation quench on the mass was similar to those in previous studies \\citep{Kauffmann2004,Haines2007,Wolf2009,Mahajan2010}, although this is the first time that galaxies of low mass ($<10^9$ M$_{\\sun}$) have been investigated in this way.  \\subsection{Comparison with galaxies with UV asymmetries}  Recently \\citet{Smith2010} investigated UV and \\Ha\\ images of the Coma cluster, and listed the galaxies with gaseous stripping events (GSE), which were detected as UV asymmetries and had colors with $NUV-i<4$. Their observed field overlapped with our field, and there were six galaxies in common (GMP2559, GMP2910, GMP3016, GMP3816, GMP4060, and GMP4232). It is striking that the UV asymmetric galaxies in our fields are only the six. This suggests that GSE galaxies have extended \\Ha\\ clouds with high probability.  The other eight parents were not listed as GSE galaxies, although all of the parents in this study satisfied the selection criteria of \\citet{Smith2010} ($i<18$, and $0.01<$z$<0.043$). As the NUV-optical color is a function of the star formation strength and the age since the last star formation \\citep[e.g.][]{Kaviraj2007}, one possibility is that the non-GSE parents are older. This can explain the red \\#3 galaxies, as age quickly makes the $g-r$ color redder. The other five star-forming galaxies and a post-starburst galaxy have young stars, and the age cannot explain the non-detection of UV asymmetry. It is possible that the UV asymmetric feature is faint because the star formation is/was weaker than other cases. It is also possible that the UV emitting stars have a symmetric distribution, while the \\Ha\\ is asymmetric. The last case suggests that the distributions of the young stars and gas differ. The sign of the different distribution was shown in our previous studies \\citep{Yagi2007,Yoshida2008}. This also suggests that the source of ionization is not always the UV photons from young stars. A detailed discussion of the source of ionization will be presented in our next paper (Yoshida et al. in preparation).  If we assume that all galaxies from which gas was removed experienced an \\Ha\\ parent phase, we can speculate on the evolution of the parents. With regard to infalling galaxies, \\citet{Graham2003} investigated two disk dwarf galaxies in the Coma cluster, which still show spiral arm patterns. The galaxies, GMP3292 and GMP3629, are thought to be transforming from spiral to passive galaxies. \\citet{Graham2003} also noted that these galaxies have large peculiar velocities relative to the Coma center. However, these two galaxies are, much redder than parents in this study ($r=16.5, g-r=0.68$ for GMP3292 and $r=18.3, g-r=0.64$ for GMP3629), but are of comparable mass (4.1$\\times 10^9$ and 8$\\times 10^8$ M$_{sun}$), with values estimated in the same way as for the parents. They might be recently evolved shapes from the parents, as they still retain the arm pattern."
    },
    "1005/1005.4665_arXiv.txt": {
        "abstract": "The {\\it Chandra\\/} Source Catalog (CSC) is a general purpose virtual X-ray astrophysics facility that provides access to a carefully selected set of generally useful quantities for individual X-ray sources, and is designed to satisfy the needs of a broad-based group of scientists, including those who may be less familiar with astronomical data analysis in the X-ray regime.  The first release of the CSC includes information about 94,676 distinct X-ray sources detected in a subset of public ACIS imaging observations from roughly the first eight years of the {\\it Chandra\\/} mission. This release of the catalog includes point and compact sources with observed spatial extents $\\lesssim 30''$.  The catalog (1)~provides access to the best estimates of the X-ray source properties for detected sources, with good scientific fidelity, and directly supports scientific analysis using the individual source data; (2)~facilitates analysis of a wide range of statistical properties for classes of X-ray sources; and (3)~provides efficient access to calibrated observational data and ancillary data products for individual X-ray sources, so that users can perform detailed further analysis using existing tools.  The catalog includes real X-ray sources detected with flux estimates that are at least 3 times their estimated $1\\,\\sigma$ uncertainties in at least one energy band, while maintaining the number of spurious sources at a level of $\\lesssim 1$ false source per field for a $100\\rm\\,ks$ observation.  For each detected source, the CSC provides commonly tabulated quantities, including source position, extent, multi-band fluxes, hardness ratios, and variability statistics, derived from the observations in which the source is detected.  In addition to these traditional catalog elements, for each X-ray source the CSC includes an extensive set of file-based data products that can be manipulated interactively, including source images, event lists, light curves, and spectra from each observation in which a source is detected. ",
        "introduction": "Ever since {\\it Uhuru\\/} \\citep{gia71}, X-ray astronomy missions have had a tradition of publishing catalogs of detected X-ray sources, and these catalogs have provided the fundamental datasets used by numerous studies aimed at characterizing the properties of the X-ray sky.  While source catalogs are the primary data products from X-ray sky surveys \\citep[e.g.,][]{gia72, for78, elv92, vog93, vog99}, the {\\it Einstein\\/} IPC catalog \\citep{har90} demonstrated the utility of catalogs of {\\it serendipitous\\/} sources identified in the fields of {\\it pointed-observation\\/} X-ray missions. More recent serendipitous source catalogs \\citep[e.g.,][]{gio90, whi94, ued05, wat08} have further expanded the list of sources with X-ray data available for further analysis by the astronomical community.  Source catalogs typically include a uniform reduction of the mission data. This provides a significant advantage for the general scientific community because it removes the need for end-users, who may be unfamiliar with the complexities of the particular mission and its instruments, to perform detailed reductions for each observation and detected source.  When compared to all previous and current X-ray missions, the {\\it Chandra\\/} X-ray Observatory \\citep[e.g.,][]{wei00, wei02} breaks the resolution barrier with a sub-arcsecond on-axis point spread function (PSF). Launched in 1999, {\\it Chandra\\/} continues to provide a unique high spatial resolution view of the X-ray sky in the energy range from $0.1$ to $10\\,\\rm keV$, over a $\\sim\\!60$--$250$ square arcminute field of view. The combination of excellent spatial resolution, a reasonable field of view, and low instrumental background translate into a high detectable-source density, with low confusion and good astrometry. {\\it Chandra\\/} includes two instruments that record images of the X-ray sky. The Advanced CCD Imaging Spectrometer \\citep[ACIS;][]{bau98, gar03} instrument incorporates ten $1024\\times1024$ pixel CCD detectors (any six of which can be active at one time) with an effective pixel size of $\\sim\\!0.5''$ on the sky, an energy resolution of order $110\\,\\rm eV$ at the Al-K edge ($1.49\\,\\rm keV$), and a typical time resolution of $\\sim\\!3.2\\,\\rm s$. The High Resolution Camera \\citep[HRC;][]{mur00} instrument consists of a pair of large format micro-channel plate detectors with a pixel size $\\sim\\!0.13''$ on the sky and a time resolution of $\\sim\\!15.6\\ \\mu\\rm s$, but with minimal energy resolution.  The wealth of information that can be extracted from identified serendipitous sources included in {\\it Chandra\\/} observations is a powerful and valuable resource for astronomy.  \\begin{figure} \\epsscale{1.0} \\plotone{f1} \\caption{\\label{fig:sources} Distribution of CSC release 1.0 master sources on the sky, in Galactic coordinates.} \\end{figure}  The aim of the \\dataset [ADS/Sa.CXO#CSC] {{\\it Chandra\\/} Source Catalog} (CSC) is to disseminate this wealth of information by characterizing the X-ray sky as seen by {\\it Chandra\\/}.  While numerous other catalogs of X-ray sources detected by {\\it Chandra\\/} may be found in the literature \\citep[e.g.,][]{zez06, bra08, rom08, luo08, mun09, elv09}, the region of the sky or set of observations that comprise these catalogs is restricted, and they are typically aimed at maximizing specific scientific goals.  In contrast, the CSC is intended to be an all-inclusive, uniformly processed dataset that can be utilized to address a wide range of scientific questions.  The CSC is intended ultimately to comprise a definitive catalog of X-ray sources detected by {\\it Chandra\\/}, and is being made available to the astronomical community in a series of increments with increasing capability over the next several years.  The first release of the CSC was published in 2009 March. This release includes information about 135,914 source detections, corresponding to 94,676 distinct X-ray sources on the sky, extracted from a {\\it subset\\/} of public imaging observations obtained using the ACIS instrument during the first eight years of the {\\it Chandra\\/} mission. The distribution of release~1 sources on the sky is presented in Figure~1.  We expect that the CSC will be a highly valuable tool for many diverse scientific investigations.  However, the catalog is constructed from pointed observations obtained using the {\\it Chandra\\/} X-ray Observatory, and is neither all-sky nor uniform in depth.  The first release of the catalog includes only point and compact sources, with observed extents $\\lesssim 30''$.  Because of the difficulties inherent in detecting highly extended sources and point and compact sources that lie close to them, and quantifying in a consistent and robust way the properties of such sources, we have chosen to exclude entire fields (or in some cases, individual ACIS CCDs) containing such sources from the first release of the CSC, as described in \\S~3.1.  Therefore, the catalog does not include sources near some of the most famous {\\it Chandra\\/} targets, and there may be selection effects that restrict the source content of the catalog and which therefore may limit scientific studies that require unbiased source samples.  The minimum flux significance threshold for a source to be included in the first release of the CSC is set conservatively, and corresponds typically to $\\sim\\!10$ detected source photons (on-axis) in the broad energy band integrated over the total exposure time.  This conservative threshold was chosen to maintain the spurious source rate at an acceptable level over the wide variety of {\\it Chandra\\/} observations that are included in this release of the catalog.  We expect to relax this criterion in future releases based on experience gained constructing the current release.  A number of other {\\it Chandra\\/} catalogs do include sources with fewer net counts than the CSC\\null.  Such fainter thresholds are attainable typically either because of specific attributes of the observations included in those catalogs, or because of the assumptions made when constructing the catalog.  As an example of the former category, the XBootes survey catalog \\citep{ken05} includes sources that are roughly a factor of two fainter than the CSC flux significance threshold.  That survey is constructed from short ($5\\rm\\,ks$) observations obtained in an area with low line-of-sight absorption.  This results in a negligible background level that substantially simplifies source detection and enables identification of sources with very few counts.  Some {\\it Chandra\\/} catalogs derived from observations with the range of exposures comparable to those that comprise the CSC \\citep[e.g.,][]{elv09, lai09, mun09} also include fainter sources.  However, in these cases the additional source fractions are in general not large, typically adding $\\lesssim 10\\%$ more sources below the CSC threshold, as described in detail in \\S~3.7.1.  For other {\\it Chandra\\/} catalogs, visual review and validation at the source level is a planned part of the processing thread \\citep[e.g.,][]{kim07, mun09}. In some cases \\citep[e.g.,][]{bro07}, visual review may be used to adjust processing parameters for individual sources.  Such manual steps are time-consuming, but enable lower significance levels to be achieved while maintaining an acceptable spurious source rate.  In contrast, the CSC catalog construction process requires that the processing pipelines run on a wide range of observations with a minimum of manual intervention.  The scope of the CSC is simply too large to require manual handling at the source level.  We do not manually inspect individual source detections, nor do we adjust source detection or processing parameters based on manual evaluation.  Instead, the CSC uses a largely automated quality assurance approach, as described in \\S~3.14.  The sky coverage of the first catalog release (Fig.~2) totals $\\sim\\!320$ square degrees, with coverage of $\\sim\\!310$ square degrees brighter than a $0.5$--$7.0\\rm\\,keV$ flux limit of $1.0\\times10^{-13}\\rm\\,erg\\,cm^{-2}\\,s^{-1}$, decreasing to $\\sim\\!135$ square degrees brighter than $1.0\\times10^{-14}\\rm\\,erg\\,cm^{-2}\\,s^{-1}$, and $\\sim\\!6$ square degrees brighter than $1.0\\times10^{-15}\\rm\\,erg\\,cm^{-2}\\,s^{-1}$. These numbers will continue to grow as the {\\it Chandra\\/} mission continues, with a 15 year prediction of the eventual sky coverage of the CSC of order 500 square degrees, or a little over 1\\% of the sky.  \\begin{figure} \\epsscale{1.0} \\plotone{f2} \\caption{\\label{fig:skycov} Sky coverage of release 1.0 of the CSC, in the ACIS broad energy band. The ordinate value is the total sky area included in the CSC that is sensitive to point sources with fluxes at least as large as the corresponding value on the abscissa.} \\end{figure}  In this paper we describe in detail the content and construction of release~1 of the CSC\\null.  However, where appropriate we also discuss in addition the steps required to process HRC instrument data used to construct release~1.1 of the catalog, since the differences in the algorithms are small.  Release~1.1 of the catalog is scheduled for spring 2010.  This paper is organized into 5 sections, including the introduction.  In \\S~2, we present a description of the catalog. This includes the catalog design goals, an outline of the general characteristics of {\\it Chandra\\/} data that are relevant to the catalog design, the organization of the data within the catalog, approaches to data access, and an outline of the data content of the catalog.  Section~3, which comprises the bulk of the paper, describes in detail the methods used to extract the various source properties that are included in the catalog, with particular detail provided when the algorithms are new or have been adapted for use with {\\it Chandra\\/} data.  A brief description of the principal statistical properties of the catalog sources is presented in \\S~4; this topic is treated comprehensively by \\citetfap.  Conclusions are presented in \\S~5.  Finally, Appendix~A contains details of the algorithm used to match source detection from multiple overlapping observations, as well as the mathematical derivation of the multivariate optimal weighting formalism used for combining source position and positional uncertainty estimates from multiple observations. ",
        "conclusions": "The {\\it Chandra\\/} Source Catalog is a general purpose virtual X-ray astrophysics facility that provides access to a carefully crafted set of scientifically useful quantities for individual X-ray sources observed by the {\\it Chandra\\/} X-ray Observatory.  The first release of the catalog was published to the astronomical community in March 2009, and includes source properties for 94,676 point and compact X-ray sources detected in a subset of public ACIS imaging observations from roughly the first eight years of the {\\it Chandra\\/} mission.  This release of the catalog includes sources with observed spatial extents are $\\lesssim 30''$, and whose flux estimates are at least 3 times their estimated $1\\,\\sigma$ uncertainties.  Observations that include substantially extended sources are not included in the first release of the catalog.  For each X-ray source, the catalog tabulates about 60 distinct measured and derived source properties, generally with associated lower and upper confidence limits, in several energy bands.  These properties are generally derived from all of the observations in which a source is detected. However, in the first catalog release, multiple observations are not {\\it combined\\/} prior to source detection, so the depth of the catalog is limited by the duration of the longest single exposure of a field.  The catalog further tabulates roughly 120 observation-specific properties for each observation of a source, again with associated lower and upper confidence limits, and in several energy bands.  Tabulated source properties include source position, spatial extent, multi-band aperture fluxes computed in several different ways, X-ray hardness ratios and spectral model fits, and intra- and inter-observation variability measures.  In addition to these ``traditional'' catalog elements, for each source detection the catalog includes an extensive set of FITS format file-based data products that can be manipulated interactively by the user, including source images, event lists, light curves, and spectra from each observation in which a source is detected.  \\begin{figure*} \\epsscale{1.0} \\plotone{f26} \\caption{\\label{fig:phot} Comparison of input ($\\rm F_0$) and measured ($\\rm F$) ACIS broad ({\\it b\\/}) band fluxes for simulated sources with power-law spectra and off-axis angles $\\theta\\leq 10'$ ({\\it left\\/}) and $\\theta > 10'$ ({\\it right\\/}).  For each bin, the horizontal line indicates the median measured flux value.  The boxes include 90\\% of the measurements in each bin, and the vertical lines indicate the extreme values.  Bins colored red include fewer than 100 measurements; bins colored blue include 100--400 measurements; bins colored black include more than 400 measurements.  The green line has a slope of 1.\\vglue 5pt} \\end{figure*}  Looking towards the future, release 1.1 of the catalog, scheduled for spring 2010, will include data from public HRC-I imaging observations and newly public ACIS imaging observations, but will otherwise retain the same limitations as release 1. In release 2, we plan to co-add multiple observations of the same field that use the same or similar instrument configurations, and that have similar spacecraft pointings (within $\\sim\\!30''$) prior to source detection, to achieve fainter limiting sensitivities in many fields.  We anticipate that new algorithms will allow this release to have a significantly fainter source detection threshold than release 1.  This release should also provide limited improvements in the area of extended source handling (for example allowing for the inclusion of exposures containing moderately extended emission from galaxy cores up to $\\sim\\!60''$ spatial scale), as well as numerous algorithm enhancements that will refine field and source property calculations."
    },
    "1005/1005.5455_arXiv.txt": {
        "abstract": "The PICASSO project is using superheated droplets of C$_4$F$_{10}$ for the direct detection of Dark Matter candidates in the {\\it spin-dependent} (SD) sector. The total setup includes 32 detectors installed in the SNOLAB underground laboratory in Sudbury (Ontario, Canada). With a concentrated effort in detector purification and with new discrimination tools now available for analysis, Picasso published competitive results in June 2009 \\cite{publi2009} and became the leading experiment in the SD sector of direct dark matter searches. The present level of sensitivity is at 0.16 pb on protons at 90\\% C.L. (M$_W$= 24GeV/c$^2$) following an analysis of two detectors only. The rest of the detectors are now in the process of being analyzed and the experimental search continues in order to further improve the limits or hopefully discover a signal of dark matter. The status of the experiment and the ongoing analysis will be presented. ",
        "introduction": "The astronomical and cosmological observations strongly suggest the presence of Dark Matter and show that only 1\\% of the matter of the Universe is luminous and 23\\% of all the matter should be of a new exotic kind \\cite{bennett} : Cold Dark Matter (CDM). The preferred candidate for particle physicists is the neutralino, the Lightest Supersymmetric Particle (LSP). It is a very Weakly Interacting Massive Particle (WIMP) and is a natural candidate in the Minimal SuperSymmetric Model (MSSM) \\cite{supersym}. PICASSO (Project In CAnada to Search for Supersymmetric Objects) is one of the many worldwide efforts to hunt for dark matter through the direct detection of neutralinos via their SD interactions with $^{19}$F nuclei \\cite{ellis91,divari,bednyako}. The goal of the PICASSO project at SNOLAB is to exploit the favorable properties of $^{19}$F by using the superheated droplet detection technique, which is based on the operation principle of the classic bubble chamber \\cite{glaser52,NC94}. ",
        "conclusions": "PICASSO recently developed a new generation of detectors using a new kind of gel without cesium chloride which is presently the main source of alpha-contamination. In this case, viscous glycerin and polyethylene glycol allow to suspend the C$_4$F$_{10}$ droplets without density matching. Many improvements have been brought to the various fabrication steps, in order to obtain the lowest possible alpha-background. The preliminary results are very encouraging and may lead to a background reduction of at least a factor 10. Presently, half of the set-up has been replaced by these new detectors and the full installation is expected to be completed before the end of the year. The analysis of the new generation of detectors is now in progress and studies of a complete discrimination between the particle induced recoils and the alpha-particles will allow PICASSO to improve substantially the level of its sensitivity in the spin-dependant sector.\\\\"
    },
    "1005/1005.2619_arXiv.txt": {
        "abstract": "The Hypervelocity Star survey presents the currently largest sample of radial velocity measurements of halo stars out to 80 kpc.  We apply spherical Jeans modeling to these data in order to derive the mass profile of the Galaxy.  We restrict the analysis to distances larger than 25 kpc from the Galactic center, where the density profile of halo stars is well approximated by a single power law with logarithmic slope between $-3.5$ and $-4.5$.  With this restriction, we also avoid the complication of modeling a flattened Galactic disk.  In the range $25 < r < 80$ kpc, the radial velocity dispersion declines remarkably little; a robust measure of its logarithmic slope is between $-0.05$ and $-0.1$.  The circular velocity profile also declines remarkably little with radius.  The allowed range of $V_c\\, (80\\, \\mathrm{kpc})$ lies between 175 and 231 $\\kms$, with the most likely value 193 $\\kms$.  Compared with the value at the solar location, the Galactic circular velocity declines by less than 20\\% over an order of magnitude in radius.  Such a flat profile requires a massive and extended dark matter halo.  The mass enclosed within 80 kpc is $6.9^{+3.0}_{-1.2} \\times 10^{11}\\ \\Msun$.  Our sample of radial velocities is large enough that the biggest uncertainty in the mass is not statistical but systematic, dominated by the density slope and anisotropy of the tracer population.  Further progress requires modeling observed datasets within realistic simulations of galaxy formation. ",
        "introduction": "\\label{sec:intro}  Measuring the mass of the Galaxy is an astonishingly difficult task. Convenient tracers of disk rotation -- stars and gas clouds -- extend only to 20 kpc \\citep[e.g.][]{sofue_etal09}.  At larger radii, statistical analysis of radial velocities must be used.  Traditional tracers at distances up to 100 kpc include globular clusters and dwarf satellite galaxies \\citep[e.g.,][]{kochanek96, wilkinson_evans99, sakamoto_etal03}.  Recently, \\citet{battaglia_etal05} derived the radial velocity dispersion profile to 120 kpc using a combined sample of globular clusters, satellite galaxies, and halo red giant stars.  They found the dispersion falling from $\\sim 120\\ \\kms$ to $\\sim 50\\ \\kms$ between 20 and 120 kpc.  In contrast, \\citet{xue_etal08} assembled a large sample of blue horizontal-branch (BHB) stars from the Sloan Digital Sky Survey Data Release 6 (SDSS DR6) and found a much flatter profile between 20 and 60 kpc.  Here we use a new spectroscopic survey \\citep{brown_etal10a} aimed at finding hypervelocity stars (HVS) to set the most precise constraint on the Galactic mass within 80 kpc.  \\citet{brown_etal10a} present a sample of 910 late B-stars and early A-stars in the Galactic halo.  Their luminosity, and therefore distance, depends on whether these stars are BHB stars or main-sequence blue stragglers with similar effective temperature and surface gravity.  The ambiguous nature of the stars is especially problematic at redder colors, $u-g > 0.6$, where the luminosity can differ by a factor of 5, as demonstrated in Fig. 4 in \\citet{brown_etal10a}.  This bimodal distribution of distance does not have a well-defined average value and therefore requires statistical sampling.  \\citet{brown_etal10a} create 100 Monte Carlo realizations of each star to sample the color and metallicity distributions and to derive the distributions of luminosity and distance.  We use this full Monte Carlo catalog of distances in our analysis, while retaining the observed values of radial velocity and its uncertainty. ",
        "conclusions": "Using maximum-likelihood analysis of a new sample of radial velocities of distant halo stars, we infer that their radial velocity dispersion profile declines little with distance from the Galactic center in the range $25 < r < 80$ kpc: $\\sigma(r) = 111 \\, (r/40\\, \\mathrm{kpc})^{-0.08}\\, \\kms$.  Spherical Jeans modeling indicates that the circular velocity profile $V_c(r)$ also falls only slightly over the same radial range and reaches between 175 and 231 $\\kms$ at 80 kpc.  The corresponding enclosed mass $M(80)$ is between $5.7\\times 10^{11}\\ \\Msun$ and $1.0\\times 10^{12}\\ \\Msun$.  A three-component model for the baryon and dark matter mass distribution gives the total virial mass of the Galaxy $M_{\\rm vir} = (1.6 \\pm 0.3)\\times 10^{12}\\ \\Msun$ at the virial radius $R_{\\rm vir} = 300$ kpc.  Our inferred mass of the Galaxy is higher than that obtained by \\citet{battaglia_etal05} ($M_{\\rm vir} \\approx 0.8\\times 10^{12}\\ \\Msun$) and \\citet{xue_etal08} ($M_{\\rm vir} \\approx 1.0\\times 10^{12}\\ \\Msun$) based on the modeling of their radial velocity datasets.  Our HVS sample contains more objects at $r > 40$ kpc than the Battaglia et al. and Xue et al. datasets.  Thus we have a stronger constraint on the shallow slope of the velocity dispersion profile and we derive a correspondingly larger mass.  Our inferred mass is consistent with the larger scale measurement by \\citet{li_white08} based on the Andromeda-Milky Way timing argument, $M_{\\rm vir} \\approx 2.4\\times 10^{12}\\ \\Msun$.  The implied dynamical mass-to-light ratio of the Galaxy, $M_{\\rm vir}/L_V \\approx 50$ in solar units, is also consistent with galaxy-galaxy weak lensing measurements by \\citet{mandelbaum_etal06}, galaxy kinematics modeling by \\citet{more_etal10}, and halo abundance matching modeling by \\citet{moster_etal10}.  Our sample of radial velocities is large enough that the biggest uncertainty in the mass estimate is not statistical but systematic. Within the framework of spherical Jeans modeling, the uncertainty is dominated by the density slope and anisotropy of the tracer population.  These parameters could be better constrained by future all-sky surveys of halo BHB stars.  Deeper surveys that target more distant stars at $r \\gtrsim 100$ kpc would be similarly dominated by uncertainty over the underlying distribution of the tracers.  The validity of spherical Jeans modeling is also limited by the presence of structure in the distribution of halo stars.  Galactic stellar halo contains remnants of disrupted satellite galaxies, some of which are still detectable as tidal streams.  Stars at $\\sim 100$ kpc from the Galactic center may not have had enough dynamical times to reach dynamical equilibrium, further limiting the application of equilibrium modeling.  A first step in the direction of circumventing these systematics was taken by \\citet{xue_etal08}, who modeled the motion of particles in realistic simulated halos.  Extension of such analysis to many different halo realizations using large samples of observed velocities may reduce the uncertainty over the global mass distribution in the Galaxy."
    },
    "1005/1005.5039_arXiv.txt": {
        "abstract": "Here we show that in the case when double peaked emission lines originate from outer parts of accretion disk, their variability could be caused by perturbations in the disk emissivity.  In order to test this hypothesis, we introduced a model of disk perturbing region in the form of a single bright spot (or flare) by a modification of the power law disk emissivity in appropriate way. The disk emission was then analyzed using numerical simulations based on ray-tracing method in Kerr metric and the corresponding simulated line profiles were obtained. We applied this model to the observed H$\\beta$ line profiles of 3C 390.3 (observed in the period 1995-1999), and estimated the parameters of both, accretion disk and perturbing region. Our results show that two large amplitude outbursts of the H$\\beta$ line observed in 3C 390.3 could be explained by successive occurrences of two bright spots on approaching side of the disk. These bright spots are either moving, originating in the inner regions of the disk and spiralling outwards by crossing small distances during the period of several years, or stationary. In both cases, their widths increase with time, indicating that they most likely decay. ",
        "introduction": "The huge amount of Active Galactic Nucleus (AGN) energy is released through accretion onto super-massive black hole (BH), supposed to exist in the center of AGN. The emission of the accretion disk is not only in the continuum, but also in the emission lines (e.g. in Fe K$\\alpha$ line) and in low ionization lines, as e.g. in broad Balmer emission lines which are seen as double peaked (DP). DP Balmer lines  are found in 20\\% of radio loud AGN  at $z < 0.4$ \\citep{eh94,eh03} and 4\\% of the Sloan digital Sky Survey (SDSS) quasars at $z < 0.33$ \\citep{st03}.  Broad, double-peaked emission lines of AGN provide dynamical evidence for presence of an accretion disk feeding a supermassive black hole in the center of AGN. But in some cases, the variability of these lines shows certain irregularities which could not be explained just by standard model of accretion disk.  The DP line profiles are often used to extract the disk parameters \\citep[see e.g.][]{ch89,ch90,eh94,eh03,st08,er09}. In a series of papers Dumont \\& Collin-Souffrin \\citep[see][and references therein]{cs87,csd90,dcs90a,dcs90b,dcs90c} investigated the radial structure and emission of the outer regions of the optically thin accretion disks in AGN and calculated detailed grid of photoionisation models in order to predict the relative strengths of low-ionization lines emitted from the disk. They obtained integrated line intensities and line profiles emitted at each radius of the disk, for its different physical parameters. They also studied the influence of the external illumination on the structure of the disk, considering the point source model, where a compact source of non-thermal radiation located at a given height illuminates the disk and the diffusion model, where the radiation of a central source is scattered back towards the disk by a hot diffusing medium. One of the first methods for calculating the profiles of optical emission lines from a relativistic accretion disk was proposed by \\citet{ch89}. The limitation to this method is that the accretion disk structure required to explain the variability of the line profiles cannot be axi-symmetric, i.e. very often the red peak is higher than blue one and that cannot be explained by this model. It is  not possible in a circular disk, in which the blue peak is always Doppler boosted to be stronger than the red peak. Therefore, \\citet{er95} adapted the circular accretion disk model to elliptical disks in order to fit the profiles of double-peaked emitters with a red peak stronger than the blue one. This model introduced eccentricity and phase angle parameters to the circular model described above, and  the pericenter distance of the elliptical orbits \\citep[see][]{er95}.  Spectroscopic monitoring of double-peaked emitters \\citep[see e.g.][]{sh01,ge07,sh09} has revealed that a ubiquitous property of the double-peaked broad emission lines is variability of their profile shapes on the timescales of months to years. DP line profiles are observed to vary on timescales of months to years, i.e. on timescales of the order of the dynamical time or longer \\citep[e.g.][]{vz91,zh91,ma93,ro98,se00,sh01,sb03,ge07}.  This slow, systematic variability of the line profile is on the timescale of dynamical changes in an accretion disk, and has been shown to be unrelated to the shorter timescale variability  seen in the overall flux in the line, due to reverberation of the variable ionizing continuum. Patterns in the variability of the broad Balmer lines are often a gradual change and reversal of the relative strengths of the blue-shifted and red-shifted peaks \\citep[see e.g.][]{ne97}.  Periodic variability of the red and blue peak strengths has also been attributed to a precessing elliptical disk, a precessing single-armed spiral (as e.g. 3C 332, 3C 390.3: \\citet{gi99}; NGC 1097: \\citet{sb03}), and a precessing warp in the disk. For instance, \\citet{wu08} computed the profiles of Balmer emission lines from a relativistic, warped accretion disk in order to explore the certain asymmetries in the double-peaked emission line profiles which cannot be explained by a circular Keplerian disk. Elliptical disks and spiral waves have been detected in cataclysmic variables \\citep{shh97,bc00}, and a radiation induced warp has been detected in the large-scale disk of the AGN NGC 4258 \\citep{mbp96}.  Spiral waves are a physically desirable model since they can be produced by instability in the vicinity of a black hole. They can play an important role in accretion disks, because they provide a mechanism for transporting angular momentum outward in the disk, allowing the gas to flow inwards, towards the central black hole. Long-term profile variability is thus a useful tool for extracting information about the structure and dynamics of the accretion disk most likely producing the double-peaked emission lines.  In this paper, we present an investigation of the disk line variations due to instability in accretion disk. First we developed a model, assuming that instability in the accretion disk affects disk emissivity. This model and some simulations of expected line profile variability are presented in \\S 2. In \\S 3 we compare the model with observations taken from long-term monitoring of 3C 390.3 \\citep{sh01} in order to obtain parameters of perturbations. In \\S 4 we discuss our results in the light of possible physical mechanisms which could cause such perturbations, and finally, in \\S 5 we outline our conclusions. ",
        "conclusions": "We developed a model of the disk perturbing region in the form of a single bright spot (or flare) by a modification of the power law disk emissivity and used this model to simulate the disk line profiles. This model has been used to fit the observed H$\\beta$ line of 3C 390.3 observed from 1995 to 1999. From this investigation we can point out the following results:  \\begin{enumerate} \\item The model which includes perturbation (bright spot) in the accretion disk can successfully explain difference in double peaked line profiles, as e.g. higher red peak even if we have the standard circular disk. The position of a bright spot has a stronger influence on one particular part of spectral line profiles (such as e.g. its core if the spot is in the central part of the disk, or \"red\" and \"blue\" wings if the spot is located on receding and approaching part, respectively). \\item Using the model for perturbing region we were able to successfully model and reproduce the observed variations of the H$\\beta$ line profile in the case of 3C 390.3, including the two large amplitude outbursts observed during the analyzed period. Therefore, the observed variations of the 3C 390.3 H$\\beta$ line could be caused by perturbations in the disk emissivity. \\item We found that two  outbursts referred by \\citet{sh01} could be explained by successive occurrences of two different bright spots on approaching side of the disk which are either moving, originating in the inner regions of the disk and spiralling outwards, or stationary. Both bright spots decay by time until they completely disappear. \\item Our results support hypothesis that the perturbations in accretion disk emissivity are probably caused by fragments in the spiral arms of the disk. \\end{enumerate}  The results presented above show that a circular disk with perturbations (bright spots) can be applied to explain different double peaked line profiles, and can be also used to trace perturbations (as well as their characteristics) from the broad double peaked line shapes."
    },
    "1005/1005.2934_arXiv.txt": {
        "abstract": "In this paper the possibility of generating large scale curvature perturbations induced from the entropic perturbations during the waterfall phase transition of standard hybrid inflation model is studied. % We show that whether or not appreciable amounts of large scale curvature perturbations  are produced during the waterfall phase transition depend crucially on the competition between the classical and the quantum mechanical back-reactions to terminate inflation. If one considers only the classical evolution of the system we show  that the highly blue-tilted entropy perturbations induce highly blue-tilted large scale curvature perturbations during the waterfall phase transition which dominate over the original adiabatic curvature perturbations.  However, we show that the quantum back-reactions of the waterfall field  inhomogeneities produced during the phase transition dominate completely over the classical  back-reactions. The cumulative quantum back-reactions of very small scales tachyonic modes terminate inflation very efficiently and shut off the curvature perturbations evolution during the waterfall phase transition. This indicates that the  standard hybrid inflation model is safe under large scale curvature perturbations during the waterfall phase transition. ",
        "introduction": "Inflation proved to be very successful both theoretically \\cite{Guth:1980zm} and observationally \\cite{Komatsu:2010fb} as a theory of early universe. The simplest models of inflation consist of a scalar field which is minimally coupled to gravity. A period of acceleration expansion is obtained if the potential is flat enough to allow for the inflaton field to slowly roll towards its minimum. With sufficient tunings in the parameters of the model, one can achieve  60 number of e-foldings or more to solve the horizon and the flatness problems of the standard  cosmology.  Hybrid inflation \\cite{Linde:1993cn, Copeland:1994vg} is an interesting model of inflation containing two scalar fields, the inflaton field and the waterfall field. In Linde's original hybrid inflation, the energy density during inflation is dominated by the vacuum while the inflaton is slowly rolling. The waterfall field is very heavy compared to the Hubble expansion rate during inflation, $H$, and it quickly rolls to its instantaneous minimum. The potential has the property that once the inflaton field reaches a critical value, $\\phi=\\phi_c$, the waterfall field becomes tachyonic triggering an instability and inflation ends quickly thereafter and the systems settles down into its global minimum.  Usually it is assumed that the waterfall field does not play any role in curvature perturbations during inflation and during phase transition. In this picture one basically borrows the technics and the results of single field inflationary models. That is, the super-horizon curvature perturbations, once they leave the Hubble radius, are frozen and remain unchanged until they re-enter the Hubble radius at a later time, such as at the time of CMB decouplings.  Here we would like to examine this picture more closely. We would like to see if hybrid inflation is safe under large scale curvature perturbations during the waterfall phase transition. If one considers only the classical evolution of the system, we show that during the phase transition  the highly-blue tilted entropy perturbations can induce large blue-tilted curvature perturbations on super-horizon scales which can completely dominate over the original adiabatic curvature perturbations. However, we show that the quantum back-reactions of the waterfall field inhomogeneities produced  during the phase transition become important before the classical back-reactions become relevant. We demonstrate that  the cumulative quantum back-reactions of the short-wavelength inhomogeneities are so strong that they uplift the tachyonic instability of the entropy perturbations and the curvature perturbations freezes.   Ideas similar to this line of thought, studying the amplifications of large scale curvature  perturbations during preheating, were studied in \\cite{Taruya:1997iv}-\\cite{Kohri:2009ac} and more recently in  \\cite{Levasseur:2010rk}.  The paper is organized as follows. In section \\ref{hybrid} we review the basics of hybrid inflation and obtain the background evolutions of the inflaton and the waterfall fields. In section \\ref{entropy} the entropy perturbations and in section \\ref{curvature} their effects on curvature perturbations are studied. In Section \\ref{back} the classical non-linear back-reactions as well as the quantum mechanical back-reactions are calculated and are compared to each other. Brief conclusions and discussions are followed in section \\ref{conclusions}.  While our work was finished the work by Lyth \\cite{Lyth:2010ch} appeared which has overlaps with our results. See also \\cite{Fonseca:2010nk} which appeared shortly after our work. ",
        "conclusions": "\\label{conclusions}  In this paper the possibility of producing large scale curvature perturbations induced from the entropy perturbations during the waterfall phase transition in hybrid inflation are studied. We have shown that whether or not appreciable amounts of large scale curvature perturbations  are produced depend crucially on the competition between classical and quantum mechanical back-reactions to terminate inflation.  If one considers only the classical back-reaction effects, one obtains a significant large scale curvature perturbations which completely dominate over the initial curvature perturbations. The induced large scale curvature perturbations would be highly blue-tilted with $n_{\\cal R} \\simeq 4$ as in \\cite{Gong:2008ni}. However, we have shown that the quantum-mechanical back-reactions of the waterfall field inhomogeneities produced during the phase transition dominate before the classical-back-reaction becomes important. In the Hartree approximation we found that the quantum back-reactions shuts off the classical tachyonic instability very efficiently  terminating inflation as well as curvature perturbations evolutions quickly after phase transition. We have shown that the main contribution to quantum back-reactions comes from the cumulation of the very small scales inhomogeneities, modes which are tachyonic during the phase transition but remain sub-horizon during entire inflationary period. We also made the interesting observation that in standard hybrid inflation where the waterfall field rapidly freezes to $\\psi=0$ at the background level, one can use $\\psi \\rightarrow  \\sqrt{ \\langle \\delta \\psi^2 \\rangle} $ as the effective classical trajectory. In summary, the quantum back-reactions have two crucial effects. First they determine the end of inflation given by Eq. (\\ref{qtoc}) and { Eq.} (\\ref{eoi}). Second, as mentioned above, they provide the effective classical trajectory via  $\\psi \\rightarrow \\sqrt{ \\langle \\delta \\psi^2 \\rangle} $ .  Although we have presented the analysis here only for the standard hybrid inflation, but we believe that this picture also holds for other models of inflation where there are sharp phase transitions at the end of inflation. This includes models of  brane inflation where inflation ends abruptly due to tachyon formation once the distance between the brane and anti-brane reaches a critical value. However, it would be interesting to see what happens in models of inflation, such as in double inflation \\cite{Silk:1986vc} where there is a mild phase transition in fields evolution during early stages of inflation. Since the difference $N_c - N_*$ in Eq. (\\ref{final}) is not very large, an appreciable amount of curvature perturbations can be created even when the quantum mechanical back-reactions are taken into account. This in turn can produce features in curvature perturbations such as in models  \\cite{Joy:2007na, Battefeld:2008py,  Battefeld:2010rf}.  It would be interesting to see the observational effects of the phase transitions during inflation as considered e.g. in \\cite{Parkinson:2004yx, Joy:2008qd}.  \\vspace{1 cm}  {\\bf"
    },
    "1005/1005.0100_arXiv.txt": {
        "abstract": "The Large Area Telescope on the {\\it Fermi} $\\gamma$-ray Space Telescope provides unprecedented sensitivity for all-sky monitoring of $\\gamma$-ray activity. It has detected a few Galactic sources, including 2 $\\gamma$-ray binaries and a microquasar. In addition, it is an adequate telescope to detect other transient sources. The observatory scans the entire sky every three hours and allows a general search for flaring activity on daily timescales. This search is conducted automatically as part of the ground processing of the data and allows a fast response to transient events, typically less than a day. Most of the outbursts detected are spatially associated with known blazars, but in several cases during the first years of observations, $\\gamma$-ray flares occurring near the Galactic plane did not reveal any initially compelling counterparts. This prompted follow-up observations in X-ray, optical, and radio to attempt to identify the origin of the emission and probe the possible existence of a class of transient $\\gamma$-ray sources in the Galaxy. Here we report on these LAT events and the results of the multiwavelength counterpart searches. ",
        "introduction": "There has been one firmly established class of variable sources in the high-energy $\\gamma$-ray sky. The Energetic $\\Gamma$-Ray Experiment Telescope ({\\it EGRET}) on the Compton $\\Gamma$-Ray Observatory discovered a population of variable $\\gamma$-ray blazars above 100 MeV\\cite{hartman:1999}. However, {\\it EGRET} also left the legacy of a large fraction of unidentified sources in the 3EG catalog. Many of these were found at low Galactic latitudes and believed to be Galactic in nature. {\\it EGRET} established $\\gamma$-ray pulsars as a Galactic population and these were thought to contribute to the unidentified sources. Several studies found indications of variability in some of the sources along the Galactic Plane (GP)\\cite{torres:2001,nolan:2003}, a behavior not expected of the pulsars, which are steady on these timescales. Additionally, no blazar counterpart was identified in several cases. This suggested the possible existence of a new Galactic $\\gamma$-ray class. Several sources showed both strong variability and a convincing lack of a blazar within the {\\it EGRET} localization errors, like 3EG\\,J0241+6103, 3EG\\,J1824-1514 and GRO\\,J1834-04.  One of these has emerged as a new type of $\\gamma$-ray source with the rediscovery of 3EG\\,J0241+6103 (COS-B 2CG\\,135+01), which was associated, although the position was uncertain, with LSI\\,$+61 \\adeg 303$, a \"prominent radio flaring star system'', high mass X-ray binary (HMXB) system at 2 kpc constituted of a B0\\,Ve star and a neutron star orbiting on an excentric orbit with a 26.5 days period. A daily/monthly variability was seen with {\\it EGRET}, but no periodicity detected\\cite{tavani:1998}. A TeV source has then been detected by MAGIC and then VERITAS, at this position, with a periodic signal modulated at the orbital period. {\\it Fermi} has detected a $\\gamma$-ray source, 0FGL\\,J0240.36113, at the position RA=40.076, DEC=61.233 with a 95\\% error radius of $1.8 \\amin$, consistent with the optical counterpart, and with a periodicity of $26.6 \\pm 0.5$\\,days, the emission peaking at the periastron, and a spectrum reminiscent of the pulsars, therefore definitely identifying in the MeV-GeV domain the first $\\gamma$-ray binary source with the HMXB system\\cite{abdo:2009ApJ701L123}.  The second interesting source is the case of 3EG\\,J1824-1514, detected by {\\it EGRET}, but without modulation, as spatially coincident with LS\\,5039, an HMXB constituted of a likely neutron star orbiting around an O6.5 star. HESS detected a periodic signal modulated at the orbital period of 3.91 days. {\\it Fermi} firmly detected at more than $12 \\sigma$ a periodic source, modulated at $3.91 \\pm 0.05$ days, in a complicated region with a very intense Galactic diffuse emission.  The most dynamic example from {\\it EGRET} is GRO\\,J1838-04 (3EG\\,J1837-0423), which produced an intense outburst in June 1995\\cite{tavani:1997}. The flux above 100\\,MeV in a 3.5 day period was found to be a factor of 7 brighter than in later observations of the region. Notably, no blazar counterpart is known within the 99\\% {\\it EGRET} error contour. The absence of a flat spectrum radio source at the levels typical of the {\\it EGRET} blazars made this a candidate for a different type of $\\gamma$-ray emitter. The proximity of such a unique outburst to the inner Galaxy led to speculation of a possible Galactic origin. The question of the progenitor of this activity remains as well as the broader question of the existence of similar sources of this type. As stated in Ref.~\\refcite{tavani:1997}, ``other unidentified {\\it EGRET} sources near the GP appear to be time variable with V$>1.5$'' in Ref.~\\refcite{mclaughlin:1996ApJ473.763}.  Finally, {\\it Fermi} has recently detected a variable high energy source coinciding with the position of the X-ray binary and microquasar Cygnus X-3, modulated at its short orbital period of 4.8 hours. Cygnus X-3 is an HMXB system located at a distance of $\\sim 7$\\,kpc, with a compact object of nature still matter in the debate, orbiting a Wolf-Rayet star\\cite{abdo:2009Science326.1512}. ",
        "conclusions": "The task of identifying counterparts for unidentified transients remains challenging. Although greatly reduced from {\\it EGRET}, the LAT error circles remain large compared to the resolution of telescopes in other wavebands. Additional arguments based on temporal and spectral characteristics are required to support firm associations. Ultimately, identifications of the Galactic transients require observations of related variability between the $\\gamma$-ray source and a candidate Galactic counterpart at lower frequency. In the absence of a detection of significant activity of the potential radio and X-ray counterparts for the LAT transients, they remain unidentified. These sources continue to be monitored regularly for $\\gamma$-ray activity as a part of the {\\it Fermi} sky survey observations."
    },
    "1005/1005.3209_arXiv.txt": {
        "abstract": "If new physics were capable to push the  neutrino-nucleon inelastic cross section three orders of magnitude beyond the standard-model (SM) prediction, then ultra-high energy (UHE) neutrinos would have already been observed at neutrino observatories. We use such a constraint to reveal information on the scale of noncommutativity (NC) $\\Lambda_{\\rm NC}$ in noncommutative gauge field theories (NCGFT) where neutrinos possess a tree-level coupling to photons in a generation-independent manner. In the energy range of interest ($10^{10}$ to $10^{11}$ GeV) the $\\theta$-expansion ($|\\theta| \\sim 1/\\Lambda_{\\rm NC}^2$) and therefore the perturbative expansion in terms of $\\Lambda_{\\rm NC}$ retains no longer its meaningful character, forcing us to resort to those NC field-theoretical frameworks involving the full $\\theta$-resummation. Our numerical analysis of the contribution to the process coming from the photon exchange, pins impeccably down a lower bound on $\\Lambda_{\\rm NC}$ to be as high as around up to 900 (450) TeV, depending on the estimates for the cosmogenic neutrino flux. If, on the other hand, one considers a surprising recent result occurred in Pierre Auger Observatory (PAO) data, that UHE cosmic rays are mainly composed of highly-ionized Fe nuclei, then our bounds get weaker, due to the diminished cosmic neutrino flux. Nevertheless, we show that even for the very high fraction of heavy nuclei in primary UHE cosmic rays, our method may still yield remarkable bounds on $\\Lambda_{\\rm NC}$, typically always above 200 TeV. Albeit, in this case one encounters a maximal value for the Fe fraction from which any useful information on $\\Lambda_{\\rm NC}$ can be drawn, delimiting  thus the applicability of our method. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.3453_arXiv.txt": {
        "abstract": "We present EIS/Hinode \\& SUMER/SoHO observations of propagating disturbances detected in coronal lines in inter-plume and plume regions of a polar coronal hole. The observation was carried out on $13^{th}$ November 2007 as JOP196/HOP045 programme. The SUMER spectroscopic observation gives the information about the fluctuation in radiance and on both resolved (Doppler shift) and unresolved (Doppler width) line-of-sight velocities whereas EIS $40\\arcsec$ wide slot images detect fluctuations only in radiance but maximizes the probability of overlapping field of view between the two instruments. From distance-time radiance maps, we detect the presence of propagating waves in a polar inter-plume region with a period of 15~min to 20~min and a propagation speed increasing from ($130~\\pm~14$)~km~s$^{-1}$\\ just above the limb, to ($330~\\pm~140$)~km~s$^{-1}$\\ around $160\\arcsec$\\ above the limb. These waves can be traced to originate from a bright region of the on-disk part of the coronal hole where the propagation speed is in the range of ($25~\\pm~1.3$)~km~s$^{-1}$ to ($38~\\pm~4.5$)~km~s$^{-1}$, with the same periodicity. These on-disk bright regions can be visualized as the base of the coronal funnels. The adjacent plume region also shows the presence of propagating disturbance with the same range of periodicity but with propagation speeds in the range of ($135~\\pm~18$)~km~s$^{-1}$ to ($165~\\pm~43$)~km~s$^{-1}$\\ only. A comparison between the distance-time radiance map of both regions, indicate that the waves within the plumes are not observable (may be getting dissipated) far off-limb whereas this is not the case in the inter-plume region. A correlation analysis was also performed to find out the time delay between the oscillations at several heights in the off-limb region, finding results consistent with those from the analysis of the distance-time maps. To our knowledge, this result provides first spectroscopic evidence of acceleration of propagating disturbances in the polar region close to the Sun (within 1.2~R/R$_{\\odot}$), which provides clues to the understanding of the origin of these waves. We suggest that the waves are likely either Alfv\\'{e}nic or fast magnetoacoustic in the inter-plume and slow magnetoacoustic in plume regions. This may lead to the conclusion that inter-plumes are preferred channel for the acceleration of the fast solar wind.\\\\ ",
        "introduction": "Coronal holes are regions of cool and low density plasma that, as such, are `dark' at coronal temperatures \\citep{1972ApJ...176..511M}. During solar minima, coronal holes are generally confined to the Sun's polar regions, while at solar maxima they can also be found at lower latitudes, usually associated with remnant active regions, as so-called `equatorial' coronal holes. The predominantly unipolar magnetic field from coronal hole regions is thought to give rise to the fast solar wind \\citep[e.g.,][]{1973SoPh...29..505K}. During solar minimum, Ulysses observations clearly show that the solar wind exhibits two modes of outflow: the fast wind, associated with polar coronal holes, with outflow speeds of $\\approx800$~km~s$^{-1}$\\ and the slow  wind with outflow speeds of $\\approx400$~km~s$^{-1}$\\ associated with equatorial regions \\citep{1997GeoRL..24.2885W,2000JGR...10510419M}. However, during solar maxima, low latitude coronal holes also show faster than average solar wind speed upto $\\approx600$~km~s$^{-1}$ \\citep{2003JGRA..108.1144Z}. Extreme-ultraviolet images of polar coronal holes reveal the presence of diffuse, spike-like or sheet-like structures called plumes \\citep{1975spre.conf..651B, 1977SoPh...53..397A}, which subtend an angle of roughly $2\\degree$ relative to Sun center at low altitude and expands super-radially with the coronal hole \\citep{1997SoPh..175..393D}. Regions between these structures are termed as inter-plumes. From VUV spectroscopy, plumes are known to be denser and cooler than the surrounding inter-plume regions \\citep[e.g.,][]{2006A&A...455..697W}, while spectral lines are observed to be broader in inter-plumes \\citep[i.e.,][]{2000SoPh..194...43B,2000ApJ...531L..79G,2003ApJ...588..566T}. However differences in mass, momentum and energy flux in plumes and in inter-plumes are still not known precisely.  There are several theoretical models which describe the role of MHD waves in the acceleration of the fast solar wind in coronal holes \\citep[see review by][and references therein]{2005SSRv..120...67O,2009LRSP....6....3C} and inter-plumes are often believed to be the primary site for this acceleration. It is further conjectured that these waves originate from the on-disk bright network regions \\citep{1997ApJ...484L..75W,2000A&A...359L...1P,2000ApJ...531L..79G,2001A&A...380L..39B}. A number of studies \\citep{1997ApJ...491L.111O, 2000ApJ...529..592O, 2001A&A...380L..39B, 2005A&A...442.1087P} have reported detection of oscillations in the off-limb regions of polar coronal holes. All of these studies point to the presence of compressional waves, thought to be slow magneto-acoustic waves \\citep{1998ApJ...501L.217D, 2006A&A...452.1059O, 2007A&A...463..713O, 2009A&A...499L..29B}. On the other hand, evidence for Alfv\\'{e}n waves propagating into the corona had been reported by \\citet{1998A&A...339..208B,2009A&A...501L..15B,2008A&A...483..271D,2009ApJ...691..794L} by studying the line width variations with height in polar coronal holes. Recent reports of detections of low-frequency ($<$ 5 mHz), propagating transverse motions in the solar corona \\citep{2007Sci...317.1192T} (from coronagraphic observation) and chromosphere \\citep{2007Sci...318.1574D} and their relationship with chromospheric spicules observed at the solar limb \\citep{2007PASJ...59S.655D} with the Solar Optical Telescope aboard Hinode \\citep{2007SoPh..243....3K} have widened interest in the subject. Recently \\citet{2009Sci...323.1582J} have reported detection of torsional Alfv\\'{e}nic motions associated with a large on-disk bright-point group. These waves are believed to be a promising candidate for the heating of the corona and acceleration of the solar wind \\citep{1971ApJ...168..509B,2005ApJ...632L..49S}.  Furthermore, it has been suggested that the fast solar wind streams originate from coronal hole funnels and are launched by reconnection at network boundaries, \\citep{2005Sci...308..519T}. Measurements of the outflow speed in the extended corona have been obtained with the Ultraviolet Coronagraph Spectrometer (UVCS) aboard SoHO \\citep[e.g.,][]{2000SoPh..197..115A,2003ApJ...588..566T,2004A&A...416..749A, 2007A&A...472..299T}. Some of these studies concluded that plumes have lower outflow speeds than inter-plume regions \\citep{1997AdSpR..20.2219N,2000ApJ...531L..79G,2000A&A...353..749W,2000A&A...359L...1P,2003ApJ...588..566T,2007ApJ...658..643R} and, hence, may not contribute significantly to the fast solar wind, whereas some other theoretical and observational studies find higher outflow speeds in plumes than in inter-plume regions for at least some altitudes above the photosphere \\citep{1999JGR...104.9947C,2003ApJ...589..623G,2005ApJ...635L.185G}. These contradictory reports led to the debate on whether plumes or inter-plumes are the preferred source regions for the acceleration of the fast solar wind. This topic is highly debated and still open for further confirmation.  Recently, \\citet{2009A&A...499L..29B} reported the detection of propagating slow magnetoacoustic waves with periods between 10~min and 30~min and speed $\\approx75$~km~s$^{-1}$ to $125$~km~s$^{-1}$\\ above the limb of a polar coronal hole. In their study, the propagating disturbances which are due to radiance perturbations are seen from the limb region up to $\\approx100\\arcsec$ above the limb. There is no discernible acceleration or deceleration of any individual feature as it propagates. In that study, the oscillations were detected in the two spectral lines of Ne~{\\sc viii}~770~\\AA\\ and Fe~{\\sc xii}~195~\\AA\\ observed with the Solar Ultraviolet Measurements Of Emitted Radiation \\citep[SUMER,][]{1995SoPh..162..189W} aboard the Solar and Heliospheric Observatory (SoHO) and with the EUV imaging spectrometer \\citep[EIS,][]{2007SoPh..243...19C} aboard Hinode, respectively.  In this paper, we combine again the capabilities of SUMER and EIS to observe the on-disk, limb and far off-limb region of the coronal hole, to search for the origin of waves close to the Sun and study their propagating nature. The plan of the paper is as follows: in section~\\ref{sec:obs}, the observations acquired for this study and the data reduction techniques are outlined. In section~\\ref{sec:result} results of the present study are presented with the distance-time radiance map analysis, power series analysis and time-delay analysis. A discussion of the observational results and a comparison with similar results are taken up in  section~\\ref{sec:discussion} and finally conclusions are drawn in section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  Analysis of Ne~{\\sc viii} and Fe~{\\sc xii} radiance x-t maps reveal the presence of outward propagating radiance disturbances in the off-limb and near off-limb region of inter-plume with periodicities of about 15~min to 20~min. From the SUMER Ne~{\\sc viii} line radiance x-t map, one can infer that the waves originate from a bright location (presumably the footpoint of a coronal funnel) and propagate towards the limb with a speed ($25~\\pm~1.3$)~km~s$^{-1}$. Around the limb the speed has increased to ($38~\\pm~4.5$)~km~s$^{-1}$, reaching ($130~\\pm~51$)~km~s$^{-1}$ off-limb. Further far off-limb, the speed of the propagation becomes ($330~\\pm~140$)~km~s$^{-1}$ as seen in the EIS Fe~{\\sc xii} line. Similar propagating disturbances are also seen in the plume region but with negligible acceleration, if any. The waves are not visible far off-limb, suggesting that they may be dissipated or, more simply, merge into the background. The waves as recorded in the inter-plume regions are either Alfv\\'{e}nic or fast magnetoacoustic in nature whereas the one seen in plumes are more likely slow magnetoacoustic type. \\citet{2005Sci...308..519T} have conjectured that the solar wind outflow is launched by reconnection at network boundaries between open flux lines and intra-network closed loops. The intra-network closed loops are pushed by supergranular convection towards the network triggering reconnection. This scenario is consistent with our identification of the origin of the propagating disturbances in inter-plume region with an on-disk bright region. These results support the view that the inter-plume regions are the preferred channel for the acceleration of the fast solar wind."
    },
    "1005/1005.3981_arXiv.txt": {
        "abstract": "The short-lived radioisotope $^{60}$Fe requires production in a core collapse supernova or AGB star immediately before its incorporation into the earliest solar system solids. Shock waves from a somewhat distant supernova, or a relatively nearby AGB star, have the right speeds to simultaneously trigger the collapse of a dense molecular cloud core and to inject shock wave material into the resulting protostar. A new set of FLASH2.5 adaptive mesh refinement hydrodynamical models shows that the injection efficiency depends sensitively on the assumed shock thickness and density. Supernova shock waves appear to be thin enough to inject the amount of shock wave material necessary to match the short-lived radioisotope abundances measured for primitive meteorites. Planetary nebula shock waves from AGB stars, however, appear to be too thick to achieve the required injection efficiencies. These models imply that a supernova pulled the trigger that led to the formation of our solar system. ",
        "introduction": "Primitive meteorites contain daughter products of the decay of short-lived radioisotopes (SLRIs) such as $^{26}$Al, $^{41}$Ca, $^{53}$Mn, and $^{60}$Fe, distributed in different minerals in a way that indicates the parent isotopes were still alive at the time of their incorporation into the refractory inclusions and chondrules that record the earliest history of the solar system. The presence of $^{60}$Fe is particularly significant, as its production requires stellar nucleosynthesis (Tachibana \\& Huss 2003; Tachibana et al. 2006). Given half-lives on the order of $\\sim 10^6$ yr, the evidence for these radioisotopes suggests that the same stellar source that synthesized them may well have triggered the collapse of the presolar dense cloud core as well, while simultaneously injecting the freshly-synthesized radioisotopes (Cameron \\& Truran 1977; Boss 1995). Supernovae resulting from massive stars in the range of $\\sim 20 M_\\odot$ to $\\sim 60 M_\\odot$ or planetary nebulae derived from intermediate-mass ($\\sim 5 M_\\odot$) AGB stars have been proposed as possible sources of all or most of these radioisotopes (e.g., Huss et al. 2009).  Shock-triggered collapse and injection into the presolar cloud (Cameron \\& Truran 1977) has been proposed and studied in detail (e.g., Boss 1995; Foster \\& Boss 1997; Vanhala \\& Boss 2002; Boss et al. 2008, 2010). Recent calculations have shown that simultaneous triggered gravitational collapse and injection of shock wave gas and dust into the collapsing cloud core is possible even when detailed heating and cooling processes in the shock-cloud interaction are included (Boss et al. 2008). Shock speeds in the range from 5 km/sec to 70 km/sec are capable of achieving simultaneous triggering and injection for a 2.2 $M_\\odot$ target cloud (Boss et al. 2010). However, these models led to considerably lower injection efficiencies than those previously estimated on the basis of models where the shock-cloud interaction was assumed to be isothermal (Boss 1995; Foster \\& Boss 1997; Vanhala \\& Boss 2002). When the injection efficiency ($f_i$) is defined to be the fraction of the incident shock wave material that is injected into the collapsing cloud core, values of $f_i \\sim 0.001$ result from the nonisothermal models (Boss et al. 2008, 2010), about 100 times lower than the values of $f_i$ found previously for strictly isothermal interactions. Considering that the shock fronts in these models contain 0.015 $M_\\odot$ of gas and dust, this means that the Boss et al. (2008, 2010) models produced nominal dilution factors $D \\sim 10^{-5}$, where $D$ is defined as the ratio of the amount of mass derived from the stellar source of the shock front that ends up in the protoplanetary disk to the amount of mass in the disk that did not derive from the stellar source. Such values appear to be much too low to explain the initial abundances inferred for typical SLRIs, which range from $\\sim 10^{-4}$ to $\\sim 3 \\times 10^{-3}$ for supernovae (Takigawa et al. 2008; Gaidos et al. 2009) and $\\sim 3 \\times 10^{-3}$ for an AGB star (Trigo-Rodr\\'iguez et al. 2009).  Boss et al. (2010) found that varying the shock speed from 5 to 70 km/sec had relatively little effect on $f_i$, while doubling the density of the target cloud could decrease $f_i$ by a factor of 3. Here we explore the effects of changes in the assumed shock wave parameters, in order to learn if higher values of $f_i$ and therefore $D$ might thereby result. In addition, we seek to learn if these shock wave variations will indicate a preference for either a supernova or an AGB star wind for triggering the formation of the solar system. ",
        "conclusions": "A new set of models with varied shock densities and thicknesses has shown that injection efficiencies $f_i$ and dilution factors $D$ can be increased by large factors ($>$ 10 and $>$ 1000, respectively), large enough to maintain the viability of this SLRI injection mechanism. Observations of supernova remnants and planetary nebulae imply that while the former shock fronts are thin enough to be suitable for SLRI injection, the latter are not. These results lend support to previous studies that have favored a supernova over an AGB star for the source of the solar system's SLRIs. Huss et al. (2009) found that an intermediate-mass AGB star could explain the production of $^{26}$Al, $^{41}$Ca, $^{60}$Fe, but not that of $^{53}$Mn. Kastner \\& Myers (1994) pointed out that AGB stars are seldom found in the vicinity of star-forming regions, so the chances of SLRI injection from a planetary nebula wind into a dense cloud core are small. The culprit appears to have been a long-forgotten supernova remnant that swept through the galaxy $\\sim$ 4.56 Gyr ago."
    },
    "1005/1005.3386_arXiv.txt": {
        "abstract": "The source HESS J1813-178 was detected in the survey of the inner Galaxy in TeV $\\gamma$-rays, and a composite supernova remnant (SNR) G12.8-0.0 was identified in the radio band to be associated with it. The pulsar wind nebula (PWN) embedded in the SNR is powered by an energetic pulsar PSR J1813-1749, which was recently discovered. Whether the TeV $\\gamma$-rays originate from the SNR shell or the PWN is uncertain now. We investigate theoretically the multiwavelength nonthermal radiation from the composite SNR G12.8-0.0. The emission from the particles accelerated in the SNR shell is calculated based on a semianalytical method to the nonlinear diffusive shock acceleration mechanism. In the model, the magnetic field is self-generated via resonant streaming instability, and the dynamical reaction of the field on the shock {\\bf is} taken into account. Based on a model which couples the dynamical and radiative evolution of a PWN in a non-radiative SNR, the dynamics and the multi-band emission of the PWN are investigated. The particles are injected with a spectrum of a relativistic Maxwellian plus a power law high-energy tail with an index of $-2.5$. Our results indicate that the radio emission from the shell can be well reproduced as synchrotron radiation of the electrons accelerated by the SNR shock; with an ISM number density of $1.4$ cm$^{-3}$ for the remnant, the $\\gamma$-ray emission from the SNR shell is insignificant, and the observed X-rays and very high energy (VHE) $\\gamma$-rays from the source are consistent with the emission produced by electrons/positrons injected in the PWN via synchrotron radiation and inverse Compton (IC) scattering, respectively; the resulting $\\gamma$-ray flux for the shell is comparable to the detected one only with a relatively larger density of about $2.8$ cm$^{-3}$. The VHE $\\gamma$-rays of HESS J1813-178 can be naturally explained to mainly originate from the nebula although the contribution of the SNR shell becomes significant with a denser ambient medium. ",
        "introduction": "\\label{sec:intro}  The VHE source HESS J1813-178 was discovered with the High Energy Stereoscopic System (HESS) in a survey of the inner Galaxy in VHE $\\gamma$-rays \\citep[][]{A05}. The VHE $\\gamma$-ray image obtained with the HESS shows a pointlike source with an extension of $\\sim2.2'$, and the observed spectrum has a hard photon index $2.09\\pm0.08$ \\citep[][]{A06}. HESS J1813-178 was also detected in VHE $\\gamma$-rays with the Major Atmospheric Gamma Imaging Cerenkov (MAGIC) telescope. The differential flux given by MAGIC between 0.4 and 10 TeV can be well described by a power law with a index of $-2.1\\pm0.2$ \\citep[][]{Al06}, which is consistent with the result obtained with the HESS.  Firstly, HESS J1813-178 was unidentified, and it was assumed to be a \"dark particle accelerator\" since no counterpart at lower frequencies was reported ever. However, the shell of the SNR G12.8-0.0 associated with HESS J1813-178 was discovered with a diameter of $\\sim2.5'$ in a new low-frequency Very Large Array (VLA) 90 cm survey \\citep[][]{Bet05}. The non-thermal radio flux densities of the shell-type SNR are $0.65\\pm0.10$ and $1.2\\pm0.08$ Jy at 20 and 90 cm, respectively \\citep[][]{Bet05}. A highly absorbed X-ray source AX J1813-178, for which the column density is $10^{23}$ cm$^{-2}$, detected with {\\it ASCA} is spatially coincident with the SNR. The X-ray emission extending to 10 keV with a sharp cutoff below 2 keV is primarily nonthermal, which can originate either from the SNR shell or form the pulsar wind nebula (PWN) inside the remnant \\citep[][]{Bet05}. Due to the high column density derived from the {\\it ASCA} data, a distance of $\\geq4$ kpc is derived for the source \\citep[][]{Bet05}.  Moreover, a soft $\\gamma$-ray source, IGR J18135-1751, was discovered as the counterpart of HESS J1813-178  \\citep[][]{Uet05}. It is persistent with a 20--100 keV luminosity of $5\\times10^{34}$ erg s$^{-1}$ for a distance of 4 kpc. \\citet[][]{Uet05} argued that the observed properties of the source in the radio and X-ray bands can be explained with the assumption that the source is a pulsar wind nebula embedded in G12.8-0.0. Furthermore, high-angular resolution X-ray observation with {\\it XMM-Newton} shows that the X-ray emitting object appears as a compact core located in the center of the radio shell-type SNR G12.8-0.0 \\citep[][]{Fet07}. They argued that the source is a composite SNR since the central object shows morphological and spectral resemblance to a PWN.  The observation with {\\it Chandra} on the SNR G12.8-0.0 indicates the X-ray source is a point surrounded by structured diffuse emission that fills the interior of the radio shell \\citep[][]{Het07}. The compact source has a spectrum characterized by a power law with an index of $\\sim1.3$, typical of young and energetic rotation-powered pulsars, and the morphology of the diffuse emission strongly resembles that of a pulsar wind nebula \\citep[][]{Het07}. Recently, an energetic pulsar PSR J1813-1749 with a period of $\\sim44.7$ ms, a characteristic age of 3.3 - 7.5 kyr, and a distance of $4.7$ kpc by assuming the association with an adjacent young stellar cluster, was discovered in a long, continuous {\\it XMM-Newton} X-ray timing observation \\citep[][]{GH09}.  The pulsar was found to be associated with the SNR G12.8-0.0, and it powers the PWN \\citep[][]{GH09}.  High-energy $\\gamma$-rays can be produced either from SNR shells in which particles are accelerated to relativistic through the first Fermi process \\citep[e.g.,][]{A05b,A07,BV06,FZ08,Fet09,MAB09}, or from PWNe powered by the pulsars inside them \\citep[e.g.,][]{Vet08, ZCF08,GSZ09}. Although the source HESS J1813-178 is pointlike in VHE $\\gamma$-rays, the possibility of the VHE $\\gamma$-rays originating from the shell of the remnant cannot be ruled out given the size of the SNR, the angular resolution of the HESS telescope, and the depth of the observations \\citep[][]{Al06}. In this paper, we study the multiband nonthermal emission from the shell of the SNR G12.8-0.0 and the nebula inside it. The emission from the particles accelerated by the SNR shock is investigated based on a semianalytical method to the nonlinear diffusive shock acceleration mechanism with a free escape boundary proposed by \\citet{Cea10fre}, in which the amplified magnetic field due to resonant streaming instability induced by cosmic rays and the dynamical feedback of this self-generated magnetic field on the shock are taken into account. On the other hand, the dynamics and the multi-band radiative properties of the PWN are investigated basically according to the model in \\citet[][]{GSZ09}, which can self-consistently describe the dynamical and radiative evolution of a pulsar wind nebula in a non-radiative supernova remnant. Recently, based on the long-term two dimensional particle-in-cell simulations, \\citet[][]{Sp08} found that the particle spectrum downstream of a relativistic shock consists of two components: a relativistic Maxwellian and a power law high-energy tail with an index of $-2.4\\pm0.1$. Different from \\citet{GSZ09}, in which a single power law injection spectrum for the electrons/positrons is employed to discuss the radiative properties during different phase of the PWN, in this paper we argue that the high-energy particles are injected with the spectrum of a relativistic Maxwellian plus a power-law high-energy tail during the evolution, and a kinetic equation is used to obtain the energy distribution of the particles. ",
        "conclusions": "\\label{sec:discussion}  The origin of the VHE $\\gamma$-rays from HESS J1813-178 is investigated in this paper. Although the source is pointlike in VHE $\\gamma$-rays, a shell origin of the  VHE $\\gamma$-rays cannot be ruled out because the PSF of the HESS telescope is very close to the size of the SNR's shell \\citep[e.g.,][]{Al06}. A SNR with a diameter of $2.5'$ was identified in the radio bands to be associated with the VHE $\\gamma$-ray source. Moreover, X-ray observations showed a PWN powered by an energetic pulsar is embedded in the SNR \\citep[][]{Bet05,Uet05,Fet07,Het07}, and the pulsar was recently discovered in a long, continuous {\\it XMM-Newton} X-ray timing observation \\citep[][]{GH09}. We apply a model which can self-consistently calculate the multiwavelength nonthermal emission both from the SNR shell and from the PWN embedded in the remnant. In the model, electrons/positrons with a spectrum of a relativistic Maxwellian plus a power-law high-energy tail are injected in the nebula during the PWN during its evolution inside the SNR; protons and electrons are accelerated by the SNR shock wave, and the spectrum of the accelerated particles are calculated with a semi-analytical non-linear model. Our results indicate that: with the parameters in Table\\ref{para}, (1) the observed emission of the shell in the radio bands can be well explained as synchrotron radiation of the electrons accelerated by the SNR shock wave, whereas the flux of p-p collisions of the accelerated protons is significantly smaller than the TeV flux observed with HESS; (2) the observed emission in the X-ray to hard X-ray band detected with {\\it XMM-Newton} and {\\it INTEGRAL}, respectively, can be explained as the synchrotron radiation of the electrons/positrons injected in the PWN; (3) the VHE $\\gamma$-rays in the TeV band are mainly produced via IC scattering of the electrons/positrons injected in the nebula on the CMB and interstellar IR photons.  With a sub-energetic explosion ($E_{\\rm sn}\\sim 0.4\\times 10^{50}$ erg), an age of $\\sim 1200$ yr, and an ambient density of $1.4$ cm$^{-3}$, the radius of the composite SNR G12.8-0.0, 1.7 pc, as well as the multiband observed nonthermal fluxes for the remnant can be reproduced within this model described in this paper. The present value of the SNR shock's velocity is about 1000 km s$^{-1}$, and the VHE $\\gamma$-rays are predominately produced from the nebula inside the remnant via IC scattering. Of course, p-p interaction from the SNR shell can be enhanced with a denser medium around the remnant (Fig.\\ref{figs6}). Even if this condition were satisfied, the $\\gamma$-rays produced via IC scattering in the PWN remain significant, hence, in this case, both the SNR shell and the PWN would contribute to the production of the observed $\\gamma$-ray flux. In GeV $\\gamma$-rays, the spectral property of the resulting emission in this SNR-dominated scenario (Fig.\\ref{figs6}) differs significantly from that in the PWN-dominated case (Fig.\\ref{g128}). The VHE $\\gamma$-rays of HESS J1813-178 detected with HESS can be naturally explained as the IC scattering of the electrons/positrons injected into the PWN although the p-p collisions become important with a denser ambient medium for the remnant. Our study give more insights on the nature of the multiband nonthermal emission of the composite SNR G12.8-0.0, even though some assumptions are made in the model."
    },
    "1005/1005.1660_arXiv.txt": {
        "abstract": "Turbulent, two-dimensional, hydrodynamic flows are characterized by the emergence of coherent, long-lived vortices without a need to invoke special initial conditions.  Vortices have the ability to sequester particles, with typical radii $\\sim 1$ mm to $\\sim 10$ cm, that are slightly decoupled from the gas.  A generic feature of discs with surface density and effective temperature profiles that are decreasing, power-law functions of radial distance is that four vortex zones exist for a fixed particle size.  In particular, two of the zones form an annulus at intermediate radial distances within which small particles reside.  Particle capture by vortices occurs on a dynamical time scale near and at the boundaries of this annulus.  As the disc ages and the particles grow via coagulation, the size of the annulus shrinks.  Older discs prefer to capture smaller particles because the gas surface density decreases with time, a phenomenon we term ``vortex aging''.  More viscous, more dust-opaque and/or less massive discs can have vortices that age faster and trap a broader range of particle sizes throughout the lifetime of the disc.  Thus, how efficiently a disc retains its mass in solids depends on the relative time scales between coagulation and vortex aging.  If vortices form in protoplanetary discs, they are important in discs with typical masses and for particles that are likely to condense out of the protostellar nebula.  Particle capture also occurs at distances relevant to planet formation.  Future infrared, submillimetre and centimetre observations of grain opacity as a function of radial distance will test the hypothesis that vortices serve as nurseries for particle growth in protoplanetary discs. ",
        "introduction": "The formation of $\\sim$ km-sized planetesimals from sub-micron-sized dust grains likely involves more than one physical process \\citep{youdin10,cy10}.  It is generally accepted that particle growth in dusty, circumstellar discs is hierarchical, eventually forming planetesimals that are the building blocks of planets \\citep{armitage07,armitage10}.  There is reasonable understanding of building small particles, but forming planetesimals from these particles remains mired in controversy.  The initial stage of growth probably proceeds through the nucleation of sub-micron-sized dust grains from the primordial nebula, which then form the monomers of fractal dust aggregates up to $\\sim 1$ mm to $\\sim 10$ cm sizes in $\\gtrsim 10^3$ yrs, beyond which growth is stalled by collisional bouncing and fragmentation \\citep{bw08,zsom10}.  In this regime, the particle dynamics and coagulation are described by Brownian motion and van der Waals forces.  One of the best astrophysical pieces of evidence for grain growth to these sizes is the detection of 3.5 cm dust emission from the classical T Tauri star TW Hya (age $\\sim 5$--10 Myr), located 56 pc away, which has a face-on circumstellar disc of radius 225 AU \\citep{wilner05}.  In standard models of protoplanetary discs, the gas pressure decreases radially outward.  Gas in the disc then moves at sub-Keplerian speeds.  Solid particles on Keplerian orbits experience a ``head wind'' with a velocity $\\sim 10^3$ cm s$^{-1}$ --- this head wind drags $\\sim 10$ cm- to $\\sim 1$ m-sized particles into the central star on time scales of $\\sim 10$--100 years \\citep{weiden77a}.  These time scales are much shorter than the characteristic time scale for these particles to collide and grow into larger particles that are unaffected by the head wind.  \\cite{safronov69} and \\cite{gw73} suggested that this difficulty can be circumvented by dust settling into the midplane of the disc and triggering gravitational instability, but \\cite{weiden80} pointed out that turbulence generated by the settling impedes the process.  Long-lived structures in the gas are a possible way of concentrating particles with sizes $\\sim 1$ mm to $\\sim 10$ cm and growing them to larger sizes.  Such structures are usually high-pressure regions in the gas that are capable of concentrating particles, which is a manifestation of Bernoulli's principle\\footnote{Bernoulli's principle states that in inviscid flows, a decrease in fluid velocity is accompanied by an increase in pressure.  Hence, locations of maximum pressure are also locations of minimum velocity.} \\citep{kundu04}.  Vortices are examples of long-lived structures --- they are spiral, non-linear motions of fluid with closed streamlines.  On Earth, ocean vortices have been observed to trap larval fish off the coast of western Australia \\citep{paterson08}.  In astrophysical settings, the possible role of vortices in planet formation was suggested by \\cite{vw46}, based on the writings of Kant, in an article entitled {\\it Die Entstehung des Planetensystems} (``The Origin of Planetary Systems'').  Since then, vortices have been suggested as possible nurseries for growth to $\\gtrsim 1$ m-sized particles \\citep{aw95,bs95,tanga96,bracco99,gl99,gl00,johansen04,bm05,fn05,kb06,inaba06,shen06,bodo07,mr09}.  In stratified protoplanetary discs, there are conceivable locations where the flow is turbulent and quasi-2D.  An attractive feature of turbulent, 2D, hydrodynamic flows is that the fluid robustly self-organizes into large, coherent, long-lived vortices amidst a backdrop of small eddies without a need to invoke special initial conditions \\citep{carn91,wm93,tabeling02}.  Such a property arises from the fact that the so-called ``vortex-stretching term'' in the vorticity equation is absent in 2D, thereby allowing an inverse cascade of energy (and forward cascade of enstrophy).  Turbulence thus becomes a friend and not a foe.  If protoplanetary discs are capable of producing turbulent, quasi-2D flows, then these may seed large-scale vortices that may survive for many orbital time scales.  Such optimism should be tempered by the fact that off-midplane vortices have not been observed in simulations of protoplanetary discs with dust settling \\citep{chiang08,johansen09}.  \\emph{The main question we are addressing in this study is: assuming vortices can be generated and sustained in discs, what sizes of particles do they capture, where are the capture locations and when does capture occurs?}  In \\S\\ref{sect:basic}, we discuss/review the order-of-magnitude physics associated with protoplanetary discs, particle-gas interactions and vortices.  In \\S\\ref{sect:static}, we start with the simpliest case of a static, minimum mass solar nebula disc and show that there are generically four vortex zones within any disc with surface density and temperature profiles that are decreasing, power-law functions of $r$.  In \\S\\ref{sect:viscous}, we consider the next level of sophistication, which is the case of an evolving, viscously-heated disc; we show that there are preferred locations and particle radii for vortex capture.  In \\S\\ref{sect:irradiated}, we generalize to the case of an evolving, viscous, irradiated disc.  We demonstrate that discs which are able to both settle particles to their midplanes and capture them via vortices have upper limits to their masses that are consistent with most observed discs.  We also show that the maximum particle radius for vortex capture in these discs is $\\sim 10$ cm, independent of disc model and weakly dependent on stellar, disc and dust properties.  In \\S\\ref{sect:discussion}, we summarize our conclusions, discuss the open questions concerning the physics of vortices and describe the relevance of our results to observations.  Table \\ref{tab:parameters} lists the fiducial values adopted for the parameters of our models. ",
        "conclusions": "\\label{sect:discussion}  \\subsection{Summary}  We have examined particle trapping by vortices in evolving, viscous and/or irradiated discs as a function of radial distance, initial disc mass, dust opacity and viscosity.  The salient points of our study are: \\begin{enumerate}  \\item If the surface density and effective temperature of protoplanetary discs are decreasing, power-law functions of the distance from the star, then all discs contain four vortex zones for a fixed particle size (Figures \\ref{fig:schematic} and \\ref{fig:mmsn}).  \\item There is an annulus at intermediate distances from the star where small particles reside and may grow via coagulation.  The size of this annulus decreases with time and also shrinks as the particles grow (Figure \\ref{fig:distances}).  Vortex capture is optimal near and at the boundaries of this annulus, and occurs within an orbital period.  \\item The optimal capture radii ($\\sim 1$ mm to $\\sim 10$ cm; Figures \\ref{fig:vis} and \\ref{fig:rad}) of particles are comparable to or smaller than the sizes of particles ($\\sim 1$ m) that drift fastest through the protoplanetary disc.  The enhanced particle concentration within the vortices may make the characteristic time for particle growth shorter than that for radial drift.  \\item Vortices in older discs prefer to capture smaller particles, a phenomenon we term ``vortex aging'' (Figures \\ref{fig:vis} and \\ref{fig:rad}).  If coagulation between sub-micron particles mixed with the gas can only produce small ($\\lesssim 1$ mm) particles, then vortices can capture them throughout the lifetime of the disc.  However, if coagulation manufactures larger ($\\sim 10$ cm) particles when the disc is young, vortices must form early to capture them.  This metaphorical dance between coagulation and vortex aging determines how efficiently vortices help the disc to retain its mass in solids.  \\item More viscous, more dust-opaque and/or less massive discs can have vortices that trap a broader range of particle sizes throughout the lifetime of the disc.  While the coagulation of grains with sizes $\\lesssim 1$ mm needs to be more in synch with the evolution of more massive discs, such discs are also expected to grow grains to larger sizes more rapidly.  \\item The maximum size of particle that can be trapped by vortices is $\\sim 10$ cm, independent of disc model and weakly dependent on stellar, disc and dust properties.  Discs where particles settle to the midplane (within $\\gtrsim 10$ yr) and are sequestered in vortices have upper limits to their masses ($\\gtrsim 0.1~M_\\odot$) that are consistent with those of most observed discs.  If vertical mixing is present (e.g., via turbulence), the maximum masses can be much smaller.  \\end{enumerate}  \\subsection{The Physics of Vortices: Open Questions}  Many open questions remain concerning the microphysics of vortices.  While we have shown that particles can be gathered by vortices, the outcome of these captures is uncertain.  Vortices may concentrate enough mass to enhance collisions rates by an order of magnitude (or more) or to trigger local gravitational instabilities \\citep{aw95,gl99,gl00,kb06}. Both paths leads to the formation of self-gravitating objects which will not drift through the disc.  However, at least some of the published simulations are run in 2D (e.g., \\citealt{davis00,inaba06,lyra09}) --- in the absence of viscosity or particles, such vortices live forever, implying that the mass of the planetesimal, embyro or planet formed depends either on the time the simulation is executed or the breakup of the vortex by the non-linear feedback of the concentrated particles.  It is more likely that if the centres of vortices are relatively quiescent, they then serve as nurseries for coagulation to occur between somewhat larger particles.  Vortex formation is also uncertain.  \\cite{lp10} show that the subcritical baroclinic instability (SBI) is a plausible way of seeding vortices.  Discs develop this \\emph{non-linear} instability when they are (radially) convectively unstable, have non-negligible thermal diffusion, and are subjected to finite vorticity perturbations ($\\sim 0.1$).  The associated Reynolds number for the shearing box simulations in this study is ${\\cal R} \\sim 10^5$; because the threshold vorticity amplitude for invoking the SBI decreases with increasing ${\\cal R}$, they speculate that the threshold amplitude could be very small (and possibly sub-sonic) in realistic discs.  Another possibility for generating vortices is via the (linear) Rossby wave instability \\citep{lovelace99,vt06}, which was invoked by \\cite{inaba06} to consider 2D vortices in protoplanetary discs.  \\cite{inaba06} found that the formation of vortices via the Rossby wave instability critically depends on the amplitude and width of an initial density bump placed within the disc.  This bump appears to be most unstable to perturbations with an azimuthal mode number of 5 (i.e., ``$m=5$'' perturbations).  \\cite{meheut10} performed 3D simulations of stratified discs and concluded that strong and persistent vortices emerge out of the flow via the Rossby wave instability.  \\cite{davis00} have noted that Rossby waves are only supported by flows with non-vanishing vorticity gradients, implying that they are relevant mostly in incompressible, inviscid flows.  The subject of vortex survivability is mired in deeper controversy.  \\cite{lp09} subject vortices embedded in a shearing sheet to 3D perturbations; only vortices with $4 \\lesssim q \\lesssim 6$ survive the elliptical instability in unstratified discs. This stability region vanishes when stratified discs are considered.  By contrast, \\cite{lithwick09} asserts that weak vortices ($q \\gg 1$) can survive in quasi-2D flows.  \\cite{lp10} find that vortices develop bursts of turbulence in their cores before surviving as weaker vortices; the SBI amplifies the vortices and the cycle restarts.  If turbulent vortex cores are generic and ubiquitous phenomena in realistic discs, then they may pose a setback to using vortices as mechanisms for concentrating particles.  Many collective properties of vortices in turbulent, 2D, hydrodynamic flows remain poorly understood.  Among these is the ``universal decay theory'', which is the empirical observation that the vortex density, radius, velocity, mean separation between vortices, enstrophy and kurtosis can be approximated by power laws of time parametrized by a single parameter --- a first-principles explanation is still being sought \\citep{tabeling02}.  Statistical theories predict vortices to always ultimately merge, in contrast to experiments that show that above a critical separation, a pair of vortices may remain separated for many dynamical times \\citep{tabeling02}.  Despite these uncertainties, vortices provide an interesting alternative to streaming instabilities \\citep[e.g.,][]{gp00, yg05, jy07, yj07, johansen09}, which require dust-to-gas ratios to approach of order unity (presumably near the disc midplane).  Vortices require no special dust-to-gas ratio to trap particles and may exist off the disc midplane.  In both cases, particles with $\\xi/\\Omega \\sim 1$ are captured and the capture size decreases as the gas in the disc dissipates.  Therefore, vortex trapping and streaming instabilities may provide complementary mechanisms for particle concentration and/or growth.  A key question to study and explore is the size distribution of planetesimals produced by each mechanism, since this might have an impact on the types of (exo)planets produced in a system.  Preliminary analyses of the initial size distribution of Solar System planetesimals suggest a broad range of values, ranging from $\\sim 1$--10 km \\citep{kb10} to $\\sim 100$--1000 km \\citep{morb09}.  \\subsection{Observational Relevance}  Our study makes falsifiable predictions about the size of particles concentrated as a function of distance from the star.  Particles smaller or larger than the optimal radii for vortex capture will be uniformly distributed throughout the disc, while those with sizes $\\sim 1$ mm to $\\sim 10$ cm will be concentrated near their respective transitional distances $r_{\\rm in}$ and $r_{\\rm out}$ (Figure \\ref{fig:schematic}).  If particle concentration leads to the growth of larger particles, these particles will then  decouple from the gas and migrate out of the vortices.  The decoupling will be stronger in the inner regions of the disc as $\\xi/\\Omega \\propto a^{-2}$ (instead of $\\propto a^{-1}$ in the outer regions).  In our Solar System, the smallest constituents (chondrules) of $\\sim 100$ km-sized asteroids have sizes $\\sim 1$ mm \\citep{hewins97}, comparable to the sizes of particles trapped by vortices.  \\cite{cuzzi08} (and references therein) have pointed out that there is a spread of about 1 Myr between the formation times of the oldest and youngest objects in the same meteorite, implying that particle growth was fairly inefficient.  This inference is in turn consistent with the limited temporal windows for particle trapping implied by vortex aging.  Another interesting property of $\\sim 10$--100 km-sized asteroids is that many of them are formed from a physically and chemically homogeneous mix of particles of a similar size, consistent with the aerodynamic sorting property of vortices.  In extrasolar settings, observations characterizing grain opacity (and hence grain growth) in protoplanetary discs can quantify particle populations as functions of radial distance, but these are still nascent and have only been accomplished for a small number of objects, e.g., HD 163296 \\citep{natta07}.  Therefore, the hypothesis that vortices serve as nurseries for particle growth will need to be tested by future infrared, submillimetre and centimetre observations of protoplanetary discs, and may lead to constraints on the time scales for grain coagulation (Figures \\ref{fig:vis} and \\ref{fig:rad}).  For example, the {\\it Atacama Large Millimeter Array (ALMA),} which operates at wavelengths between 0.3 and 3.6 mm, might be able to detect disc features associated with vortex capture.  Coagulation-fragmentation simulations can subsequently be performed to convert the observed fluxes and $\\alpha$ values into constraints on the grain properties (e.g., mass/size distribution, porosity; \\citealt{birn10}).  Even if vortices do not survive long enough to create self-gravitating structures, they will certainly play a strong role in redistributing matter throughout the disc.  Whether the redistribution of matter by the vortices plays any significant role in the eventual formation of planetesimals is unknown.  An implication of a size-dependent redistribution of matter is that if the amount of electric charge carried by a particle is proportional to its size, then charge separation in protoplanetary discs may be a fairly common phenomenon.  Our study has shown that if vortices form in protoplanetary discs, they are important in discs with typical masses and for particles that are likely to condense out of the protostellar nebula.  The capture of particles also occurs at distances relevant to planet formation.  With this study, we hope to (re)ignite the debate in connecting the microphysics of vortices with the global properties of protoplanetary discs, as the first step towards understanding the \\emph{efficiency} of planetesimal --- and eventually planet --- formation."
    },
    "1005/1005.0385_arXiv.txt": {
        "abstract": "Bulges are commonly believed to form in the dynamical violence of galaxy collisions and mergers. Here we model the stellar kinematics of the Bulge Radial Velocity Assay ({\\sl BRAVA}), and find no sign that the Milky Way contains a classical bulge formed by scrambling pre-existing disks of stars in major mergers. Rather, the bulge appears to be a bar, seen somewhat end-on, as hinted from its asymmetric boxy shape.  We construct a simple but realistic $N$-body model of the Galaxy that self-consistently develops a bar. The bar immediately buckles and thickens in the vertical direction. As seen from the Sun, the result resembles the boxy bulge of our Galaxy. The model fits the {\\sl BRAVA} stellar kinematic data covering the whole bulge strikingly well with no need for a merger-made classical bulge. The bar in our best fit model has a half-length of $\\sim 4\\;\\kpc$ and extends $20^\\circ$ from the Sun-Galactic Center line. We use the new kinematic constraints to show that any classical bulge contribution cannot be larger than $\\sim$ 8\\% of the disk mass. Thus the Galactic bulge is a part of the disk and not a separate component made in a prior merger. Giant, pure-disk galaxies like our own present a major challenge to the standard picture in which galaxy formation is dominated by hierarchical clustering and galaxy mergers. ",
        "introduction": "\\label{sec:intro}  Astronomers commonly make the Copernican assumption that our Milky Way is in no way unusual; then they can exploit the fact that we live in it to study galaxy formation in special detail. Past assumptions about our Galactic bulge grew out of our developing understanding of galaxy formation.  It is well known that spiral galaxies consist of three main components, an invisible dark matter halo, an embedded, flat disk, and a central bulge.  The bulge of our Galaxy is $>$99\\% made of stars that are at least 5 Gyr old \\citep{cla_etal_08} with a wide range of metal abundances \\citep{mcw_ric_94,ful_etal_06,zoc_etal_08}.  In this respect and many others, big bulges are similar to (diskless) elliptical galaxies. The formation of ellipticals is well understood.  Hierarchical gravitational clustering of initial fluctuations in the cosmological density results in galaxy collisions and mergers that scramble flat disks into rounder ellipticals \\citep{toomre_77_merger,whi_ree_79,ste_nav_02,nak_nom_03}.  Significant energy has been invested in developing this very successful theory of galaxy formation, and it was natural to think that our Galactic bulge is a product of it.  There is little danger that the picture is fundamentally wrong \\citep{binney_04_sydney}.  But it is incomplete.  The theme of this paper is that our Galactic bulge is indeed normal but that it is prototypical of different formation processes than are usually assumed.  A complementary suite of evolution processes shapes isolated galaxies.  They evolve by rearranging energy and angular momentum; this grows central components that masquerade as classical bulges but that formed directly out of disks without any collisions \\citep{kormen_93,kor_ken_04}.  To distinguish them from merger remnants, we call them ``pseudobulges''.  They come in two varieties.  Some are flattened; they are grown out of disk gas transported inward by nonaxisymmetries such as bars.  Another variety is recognized only in edge-on galaxies.  When bars form out of disks, they buckle vertically and heat themselves into thickened structures that look box-shaped when seen edge-on \\citep{com_san_81,com_etal_90,rah_etal_91}.  Our Galaxy contains such a box-shaped \\citep{mai_etal_78,wei_etal_94,dwe_etal_95} pseudobulge \\citep{kor_ken_04}.  The identification of this boxy structure as an edge-on bar is particularly compelling because infrared imagery shows a parallelogram-shaped distortion \\citep{mai_etal_78,wei_etal_94,dwe_etal_95} that is naturally explained as a perspective effect: the near end of the bar is closer to us than the far end. So its vertical extent is taller on the near side than on the far side \\citep{bli_spe_91}. \\citet{zhao_96} developed the first rapidly rotating bar model that fitted this distortion.  Zhao's model was based on the \\citet{schwar_79} orbit superposition technique, so it was self-consistent and in steady state, but it did not evolve into that state from plausible initial conditions.  Also, little stellar kinematic data were available to constrain Zhao's steady-state model and early $N$-body models \\citep{fux_97,fux_99,sev_etal_99}, and subsequent radial velocity data from a survey of planetary nebulae, although compared with a range of dynamical models \\citep{bea_etal_00}, led to only limited conclusions because of the small numbers and uncertain population membership of the planetary nebulae.  In this paper, we simulate numerically the self-consistent formation of a bar that buckles naturally into a thickened state, and we scale that model to fit new kinematic data on bulge rotation and random velocities.  The radial velocity observations are provided by the Bulge Radial Velocity Assay ({\\sl BRAVA}; \\citealt{ric_etal_07,how_etal_08}).  This is a spectroscopic survey of the stellar radial velocities of M-type giant stars whose population membership in the bulge is well established.  These giants provide most of the 2 $\\mu$m radiation whose box-shaped light distribution motivates bar models.  {\\sl BRAVA} emphasizes measurements in two strips at latitude $b=-4^\\circ$ and $b=-8^\\circ$ and at longitude $-10^\\circ$ $< l <$ $+10^\\circ$.  A strip along the minor axis ($l \\equiv 0^\\circ$) has also been observed.  We use nearly 5,000 stellar radial velocities in this report.  A preliminary analysis of data found strong cylindrical rotation \\citep{how_etal_09} consistent with an edge-on, bar-like pseudobulge, although a precise fit of a bar model to the data was not available. This success leads us here to construct a full evolutionary $N$-body model that we can fit to the radial velocity data.    \\begin{figure} \\centerline{ \\includegraphics[angle=0.,width=\\hsize]{cont.ps} } \\caption{Upper three panels: Face-on and side-on views of the surface density of our best-fitting model as seen from far away.  The Sun's position 8.5~kpc from the Galactic center is marked along the $+x$ axis.  The Galaxy rotates clockwise as seen in the face-on projection. Bottom panel: Model surface brightness map in Galactic coordinates as seen from the Sun's location.  Our perspective makes the box-shaped, edge-on bar look taller on its nearer side.  The Galactic boxy bulge is observed to be similarly distorted.} \\label{fig:cont} \\end{figure} ",
        "conclusions": "In our models, a bar develops self-consistently from the initially unbarred, thin disk.  Bar formation enhances the radial streaming motions of disk particles, so the radial velocity dispersion quickly grows much bigger than the vertical one.  Consequently the disk buckles vertically out of the plane like a fire hose; this is the well known buckling or corrugation instability \\citep{toomre_66,com_etal_90,rah_etal_91}.  It raises the vertical velocity dispersion and increases the bar's thickness.  This happens on a short dynamical timescale and saturates in a few hundred million years.  The central part of the buckled bar is elevated well above the disk mid-plane and resembles the peanut morphology of many bulges including the one in our Galaxy \\citep{com_etal_90,rah_etal_91}.  Out of a large set of $N$-body models, we find the one that best matches our {\\sl BRAVA} kinematic data after suitable mass scaling. The barred disk evolved from a thin exponential disk that contains $M_{\\rm d}=4.25\\times 10^{10}\\Ms$, about 55 \\% of the total mass at the truncation radius (5 scale-lengths). The scale-length and scale-height of the initial disk are $\\sim$ 1.9~kpc and 0.2~kpc, respectively. The disk is rotationally supported and has a Toomre-Q of 1.2.  The amplitude of the final bar is intermediate between the weakest and strongest bars observed in galaxies. The bar's minor-to-major axial ratio is about 0.5 to 0.6, and its half-length is $\\sim$ 4 kpc. Figure~1 (top three panels) shows face-on and side-on views of the projected density of the best-fitting model. A distinctly peanut shaped bulge is apparent in the edge-on projection.  Figure~1 (bottom panel) shows the surface brightness distribution in Galactic coordinates as seen from the Sun's vantage point.  Nearby disk stars dilute the peanut shape, but the bar still looks boxy.  Moreover, from close up, an asymmetry in the longitudinal direction is apparent; this means that the bar cannot be aligned with the direction from the Sun to the Galactic center.  Rather, its near end is at positive Galactic longitude, so it looks taller in that quadrant, and it extends farther from the Galactic center on the near side than on the far side.  Both the boxy shape and the asymmetry are in good agreement with the morphology revealed by the COBE satellite near-infrared images \\citep{wei_etal_94,dwe_etal_95}.  Figure~2 compares the best-fitting model kinematics (solid lines) with the mean velocity and velocity dispersion data from the {\\sl BRAVA} and other surveys \\citep{ran_etal_09}.  All velocities presented here have been converted to Galactocentric values (the line-of-sight velocity that would be observed by a stationary observer at the Sun's position).  For the first time, our model is able simultaneously to match the mean velocities and velocity dispersions along two Galactic latitudes ($-4^\\circ$ and $-8^\\circ$) and along the minor axis.  Figure~3 constrains the angle between the bar and the line that connects the Sun to the Galactic center.  It compares the model results with the data in the $b=-4^\\circ$ major-axis strip as we vary the above angle.  Clearly the smallest bar angles give the best match to the velocity dispersions.  Intriguingly, we find that the velocity dispersions provide much stronger constraints than the mean velocity profile.  A bar angle of $0^\\circ$ also matches the kinematics well. However, the photometric asymmetry excludes a bar that is pointed at the Sun.  We therefore conclude that the overall best-fitting model has a bar angle of $\\sim 20^\\circ$.  Other studies converged on a similar bar angle \\citep{sta_etal_97,freude_98,fux_97,fux_99,bis_ger_02}.  The excellent match to the data in Figures 1 -- 3 strongly supports the suggestion that the boxy pseudobulge of the Milky Way is an edge-on, buckled bar that evolved from a cold, massive disk. The thickened disk in the pseudobulge-forming process may have contributed to the thick disk of the Milk Way, as hinted from chemical similarities of Galactic bulge and local thick disk stars \\citep{alv_etal_10,ben_etal_10}.  The model in Figures 1 -- 3 contains no classical bulge component.  Could a small classical bulge also be present?  Could it have been spun up by the formation of a bar, flattened thereby and made hard to detect?  To constrain such multi-component models with {\\sl BRAVA} kinematics, we also constructed models with a pre-existing classical bulge.  The distribution function for the live classical bulge component was generated iteratively \\citep{deb_sel_00} to ensure that both the disk and the classical bulge were initially in equilibrium, and we required that the bulge parameters are close to the fundamental plane for classical bulges and ellipticals \\citep{kor_etal_09}. The setup of the disk is the same as in the disk-only model.  We show in Figure~4 that inclusion of a classical bulge -- one widely thought to be typical of Sbc spiral galaxies like our own -- greatly worsens the model fit to the data.  The degradation is especially obvious along the Galaxy's minor axis. Including a classical bulge with just 8\\% of the disk mass considerably worsens the fit of the model to the data. Our models rule out that the Milky Way has a significant classical bulge whose mass is $>\\sim$ 15 \\% of the disk mass.   \\begin{figure*}[!ht] \\centerline{ \\includegraphics[angle=0.,width=0.8\\hsize]{solar_v_3panel_vary_classical_bulge2.ps}} \\caption{Fits to the kinematic data (cf.~Figure~2) of models that include a pre-existing classical bulge. The heavy black lines from Figure 2 represent the model without a classical bulge.  The red, green, and blue lines are for models whose classical bulges have masses of 8\\%, 15\\%, and 30\\%, respectively, of the disk mass $M_{\\rm disk}$.  Including a classical bulge significantly worsens the model fits to the data, especially along the minor axis. } \\label{fig:varybulge} \\end{figure*}  Could a smaller, merger-built bulge hide inside the boxy bar?  The only result that we are aware of that might point to such a conclusion is the observed drop in stellar metal abundances with increasing height above the Galactic plane \\citep{zoc_etal_08,zoccal_10}. \\citet{zoc_etal_08} argue that this means that the bulge must consist of both a classical and an edge-on bar component. However, no kinematic gradient or transition corresponding to the abundance gradient is observed.  Moreover, an abundance gradient can be produced within the context of secular pseudobulge formation if some of the vertical thickening is produced by resonant heating of stars that scatter off the bar \\citep{pfe_nor_90}.  If the most metal-poor stars are also the oldest stars, then they have been scattered for the longest time and now reach the greatest heights.   Our results have important implications for galaxy formation.  We demonstrate that the boxy pseudobulge is not a separate component of the Galaxy but rather is an edge-on bar.  Bars are parts of disks.  To be sure, the stars in our Galactic bar are older than most disk stars. But those stars could have formed over a short period of time but long before the bar structure formed \\citep{wyse_99,freema_08_IAU}, their old age \\citep{zoc_etal_03,ful_etal_07} is therefore not an argument against the internal secular evolution model. Our kinematic observations show no sign that the Galaxy contains a significant merger-made, ``classical'' bulge.  So, from a galaxy formation point of view, we live in a pure-disk galaxy.  Our Galaxy is not unusual: it is very similar to another giant edge-on galaxy with a boxy bulge, NGC~4565.  \\citet{kor_bar_10} recently show that NGC~4565 does not contain even a small classical bulge component and that it therefore is another giant, pure-disk galaxy that contains no sign of a merger remnant.  In fact, giant, pure-disk galaxies are common in environments like our own that are far from rich clusters of galaxies \\citep{kor_etal_10}. Classical-bulge-less, pure-disk galaxies present an acute challenge to the current picture of galaxy formation in a Universe dominated by cold dark matter -- growing a giant galaxy via hierarchical clustering ($\\Vc \\simeq 220$ km~s$^{-1}$ in the Milky Way) involves so many mergers that it seems almost impossible to avoid forming a substantial classical bulge \\citep{pee_nus_10,age_etal_10}.  How did our Galaxy grow so large with no observational sign that it suffered a major merger after the time 9 -- 10 Gyr ago \\citep[e.g.][]{win_kep_08} when the first disk stars formed?"
    },
    "1005/1005.2441_arXiv.txt": {
        "abstract": "We determine the phase diagram for dense carbon/ oxygen mixtures in White Dwarf (WD) star interiors using molecular dynamics simulations involving liquid and solid phases.   Our phase diagram agrees well with predictions from Ogata et al. and Medin and Cumming and gives lower melting temperatures than Segretain et al.  Observations of WD crystallization in the globular cluster NGC 6397 by Winget et al. suggest that the melting temperature of WD cores is close to that for pure carbon.  If this is true, our phase diagram implies that the central oxygen abundance in these stars is less than about 60\\%.  This constraint, along with assumptions about convection in stellar evolution models, limits the effective $S$ factor for the $^{12}$C($\\alpha,\\gamma$)$^{16}$O reaction to $S_{300}\\leq170$ keV barns. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.0996_arXiv.txt": {
        "abstract": "The Extragalactic Background Light (EBL) includes photons with wavelengths from ultraviolet to infrared, which are effective at attenuating gamma rays with energy above $\\sim 10$ GeV during propagation from sources at cosmological distances. This results in a redshift- and energy-dependent attenuation of the $\\g$-ray flux of extragalactic sources such as blazars and Gamma-Ray Bursts (GRBs).  The Large Area Telescope onboard {\\em Fermi} detects a sample of $\\g$-ray blazars with redshift up to $z\\sim 3$, and GRBs with redshift up to $z\\sim 4.3$. Using photons above 10 GeV collected by {\\em Fermi} over more than one year of observations for these sources, we investigate the effect of $\\g$-ray flux attenuation by the EBL. We place upper limits on the $\\g$-ray opacity of the Universe at various energies and redshifts, and compare this with predictions from well-known EBL models.  We find that an EBL intensity in the optical--ultraviolet wavelengths as great as predicted by the ``baseline'' model of \\citet{Stecker06} can be ruled out with high confidence. ",
        "introduction": "The {\\em Fermi} Gamma Ray Space Telescope was launched 2008 June 11, to provide an unprecedented view of the $\\g$-ray Universe. The main instrument onboard {\\em Fermi}, the Large Area Telescope (LAT), offers a broader bandpass \\citep[$\\sim$ 20~MeV to over 300~GeV;][]{FERMI} and improved sensitivity (by greater than an order of magnitude) than that of its predecessor instrument EGRET onboard the {\\em Compton Gamma Ray Observatory}~\\citep{EGRET}, and the Italian Space Agency satellite {\\em AGILE}~\\citep{AGILE}, which was launched in 2007.  The LAT observes the full sky every 3 hr in survey mode leading to a broadly uniform exposure with less than $\\sim 15\\%$ variation. The Gamma-ray Burst Monitor, the lower energy ($\\sim 8$~keV -- 40~MeV) instrument onboard {\\em Fermi}, observes the full un-occulted sky at all times and provides alerts for transient sources such as GRBs.  A major science goal of {\\em Fermi} is to probe the opacity of the Universe to high-energy (HE) $\\g$-rays as they propagate from their sources to Earth. Such energetic photons are subject to absorption by production of electron-positron ($e^-e^+$) pairs while interacting with low energy cosmic background photons~\\citep{Nishikov61,Gould66,Fazio70} if above the interaction threshold: $\\epsilon_{\\rm thr}=(2 m_e c^2)^2/(2E(1-\\mu))$ where $\\epsilon$ and $E$ denote the energies of the background photon and $\\g$ ray, respectively, in the comoving frame of the interaction, $m_ec^2$ is the rest mass electron energy, and $\\theta = \\arccos(\\mu)$ the interaction angle.  Because of the sharply peaked cross section close to threshold, most interactions are centered around $\\epsilon^*\\approx 0.8(E/{\\rm TeV})^{-1}$eV for a smooth broadband spectrum.  Thus, the extragalactic background light (EBL) at UV through optical wavelengths constitutes the main source of opacity for $\\g$-rays from extragalactic sources (Active Galactic Nuclei: AGN and GRBs) in the LAT energy range.  The effect of absorption of HE $\\g$-rays is then reflected in an energy- and redshift dependent softening of the observed spectrum from a distant $\\g$-ray source. The observation, or absence, of such spectral features at HEs, for a source at redshift $z$ can be used to constrain the $\\g\\g\\to e^+e^-$ pair production optical depth, $\\tau_{\\g\\g}(E,z)$.  The EBL is dominated by radiation from stars, directly from their surface and via reprocessing by dust in their host galaxies, that accumulated over cosmological evolution. Knowledge of its intensity variation with time would probe models of galaxy and star formation.  The intensity of the EBL from the near-IR to ultraviolet is thought to be dominated by direct starlight emission out to large redshifts, and to a lesser extent by optically bright AGN.  At longer wavelengths the infrared background is produced by thermal radiation from dust which is heated by starlight, and also emission from polycyclic aromatic hydrocarbons~\\citep[see e.g.][]{driver08}.  Direct measurements of the EBL is difficult due to contamination by foreground zodiacal and Galactic light~\\citep[e.g.,][]{Hauser01}, and galaxy counts result in a lower limit since the number of unresolved sources is unknown~\\citep[e.g.,][]{Madau00}. Furthermore, evolution of the EBL density in the past epochs ($z>0$) that is required to calculate the $\\g$-ray flux attenuation from distant sources cannot be addressed by measuring the EBL density at the present epoch (z=0). Hence, several approaches have been developed to calculate the EBL density as a function of redshift. The models encompass different degrees of complexity, observational constraints and data inputs. Unfortunately, the available direct EBL measurements do not constrain these models strongly at optical-UV wavelengths due to the large scatter in the data points. A description of the different models is beyond the scope of this work; we refer the reader to the original works on the various EBL models \\citep[e.g.,][]{Salamon98, Stecker06, Kneiske02, Kneiske04, Primack05, Gilmore09, Franceschini08, Razzaque09, Finke09_model}. We note that all recent EBL models, and in particular all models used in this paper, use almost identical parameters of a $\\Lambda$CDM cosmology model.     For the analyses presented in this work we have made use of the optical depth values $\\tau(E,z)$ provided by the authors of these EBL models. These models are available via webpages \\footnote{ \\url{http://www.physics.adelaide.edu.au/~tkneiske/Tau_data.html} for Kneiske 2004; \\url{http://www.phy.ohiou.edu/~finke/EBL/index.html} for Finke et al. 2010}, analytical approximations (as in, e.g., \\citet{Stecker06}), published tables (as in, e.g., \\citet{Franceschini08}) or via private communications (which is the case for, e.g., \\citet{Salamon98, Primack05, Gilmore09, Finke09_model} for this work). Since the optical depth values are usually available in tabular form, for exact values of observed energy $E$ and redshift $z$ a linear interpolation of $\\tau(E,z)$ is used for arbitrary values of $E$ and $z$ in our calculations below.   The range of predictions by these EBL models is illustrated in Figure~\\ref{fig:tau_vs_energy} as a function of observed $\\g$-ray energy for sources at different redshifts. The Universe is optically thin ($\\tau_{\\g\\g} < 1$) to $\\gamma$-rays with energy below $\\simeq 10$ GeV up to redshift $z\\simeq 3$, independently of the model (see also \\citet{dieter}). This is due to the rapid extinction of EBL photons shortwards of the Lyman limit. Gamma rays below $\\sim 10$~GeV are not attenuated substantially because of faint far-UV and X-ray diffuse backgrounds.   \\begin{figure} \\epsscale{0.9} \\plotone{f1} \\caption{Attenuation as a function of observed gamma-ray energy for the EBL models of \\citep {Franceschini08} and \\citep{Stecker06}. These models predict the minimum and maximum absorption of all models in the literature, and thus illustrate the range of optical depths predicted in the {\\em Fermi}-LAT energy range.} \\label{fig:tau_vs_energy} \\end{figure}  The primary sources of HE extragalactic $\\g$-rays are blazars and gamma ray bursts (GRBs). Blazars are active galactic nuclei (AGN) with relativistic plasma outflows (jets) directed along our line of sight. GRBs are associated with the core collapse of massive stars, or might be caused by binary mergers of neutron stars or neutron star - black hole systems. Some GRBs produce beamed high-energy radiation similar to the case of blazars but lasting for a short period of time. GRBs have not been used to constrain EBL absorption during the pre-{\\em Fermi} era mainly because of a lack of sensitivity to transient objects above 10 GeV. The sensitivity of EGRET decreased significantly above 10 GeV, and the field-of-view (FoV) of TeV instruments is small (typically $2-4^\\circ$) to catch the prompt phase where most of the HE emission occurs. The new energy window ($10- 300$ GeV) accessible by {\\em Fermi}, and the wide FoV of the LAT, makes GRBs interesting targets to constrain EBL absorption in this energy band.  Evaluating the ratio of the putatively absorbed to unabsorbed fluxes from a large number of distant blazars and GRBs observed by {\\em Fermi} could result in interesting EBL constraints, as proposed by \\citet{chen04}, although intrinsic spectral curvature \\citep[e.g.,][]{Massaro06} or redshift dependent source internal absorption \\citep{Reimer07} could make this, or similar techniques, less effective. \\citet{georgan08} have proposed that Compton scattering of the EBL by the radio lobes of nearby radio galaxies such as Fornax A could be detectable by the {\\em Fermi}-LAT.  If identified as unambiguously originating from such process, a LAT detection of Fornax A could constrain the local EBL intensity.  Because the e-folding cutoff energy, $E(\\tau_{\\g\\g}=1)$, from $\\g\\g$ pair production in $\\g$-ray source spectra decreases with redshift, modern Cherenkov $\\g$-ray telescopes are limited to probing EBL absorption at low redshift due to their detection energy thresholds typically at or below 50~GeV to 100~GeV \\citep{Hinton09}. Ground-based $\\g$-ray telescopes have detected 35 extragalactic sources to date\\footnote{e.g.,\\url{http://www.mpi-hd.mpg.de/hfm/HESS/pages/home/sources/}, \\url{http://www.mppmu.mpg.de/~rwagner/sources/}}, mostly of the high-synchrotron peaked (HSP) BL Lacertae objects type.  The most distant sources seen from the ground with a confirmed redshift are the flat spectrum radio quasar (FSRQ) 3C~279 at $z=0.536$ \\citep{Albert08} and PKS~1510-089 at $z=0.36$ \\citep{PKS1510}.  Observations of the closest sources at multi-TeV energies have been effective in placing limits on the local EBL at mid-IR wavelengths, while spectra of more distant sources generally do not extend above 1 TeV, and therefore probe the optical and near-IR starlight peak of the intervening EBL \\citep[e.g.,][]{Stecker93, Stanev98, Schroedter05,Aharonian99,Aharonian02,Costamante04, Aharonian06a, Mazin07, Albert08, Krennrich08, Finke09}.  The starting point for constraining the EBL intensity from observations of TeV $\\g$-rays from distant blazars with atmospheric Cherenkov telescopes is the assumption of a reasonable intrinsic blazar spectrum, which, in the case of a power law, $dN/dE \\propto E^{-\\Gamma_{int}}$ for example, that is not harder than a pre-specified minimum value, e.g., $\\Gamma_{int}\\geq\\Gamma_{min}=0.67$ or 1.5. Upper limits on the EBL intensity are obtained when the reconstructed intrinsic spectral index from the observed spectrum, $\\Gamma_{obs}$, presumably softened by EBL absorption of very high energy (VHE) $\\g$-rays, is required to not fall below $\\Gamma_{int}$.  The minimum value of $\\G$ has been a matter of much debate, being reasoned to be $\\G_{int}=1.5$ by \\citet{Aharonian06a} from simple shock acceleration theory and from the observed spectral energy distribution (SED) properties of blazars, while \\citet{stecker07} argued for harder values (less than 1.5) under specific conditions based on more detailed shock acceleration simulations. \\citet{Katarzynski06} suggested that a spectral index as hard as $\\Gamma_{int}=0.67$ was possible in a single-zone leptonic model if the underlying electron spectrum responsible for inverse-Compton emission had a sharp lower-energy cutoff. \\citet{boett08} noted that Compton scattering of the cosmic microwave background radiation by extended jets could lead to harder observed VHE $\\g$-ray spectra, and \\citet{aharonian08} have argued that internal absorption could, in some cases, lead to harder spectra in the TeV range as well.  A less model dependent approach uses the (unabsorbed) photon index as measured in the sub-GeV range as the intrinsic spectral slope at GeV-TeV energies. This method has recently been applied to PG~1553+113 \\citep{lat1553} and 1ES~1424+240 \\citep{lat1424,prandini10} to derive upper limits on their uncertain redshifts, and to search for EBL-induced spectral softening in {\\em Fermi} observations of a sample of TeV-selected AGN \\citep{TeVselected}.  Attenuation in the spectra of higher redshift objects ($z \\gtrsim 1$) may be detectable at the lower energies that are accessible to the {\\em Fermi}-LAT, i.e., at $E\\approx 10-300$ GeV. Gamma rays at these energies are attenuated mainly by the evolving UV background, which is produced primarily by young stellar populations and closely traces the global star-formation rate. Observations with {\\em Fermi} of sources out to high redshift could therefore reveal information about the star-formation history of the Universe, as well as the uncertain attenuation of UV starlight by dust.  In this paper we present constraints on the EBL intensity of the Universe derived from {\\em Fermi}-LAT observations of blazars and GRBs. The highest-energy $\\g$-rays from high redshift sources are the most effective probe of the EBL intensity, and consequently a powerful tool for investigating possible signatures of EBL absorption.  In contrast to ground-based $\\g$-ray detectors, {\\em Fermi} offers the possibility of probing the EBL at high redshifts by the detection of AGN at $\\gtrsim 10$ GeV energies out to $z>3$, and additionally by the detection of GRB 080916C at a redshift of $\\sim 4.35$ \\citep{abdo09,Greiner09}. GRBs are known to exist at even higher redshifts \\citep[GRB 090423 is the current record holder with $z\\sim$ 8.2]{Tanvir09}. Therefore observations of these sources with {\\em Fermi} are promising candidates for probing the optical-UV EBL at high redshifts that are not currently accessible to ground-based (Cherenkov) telescopes.  In Section~\\ref{sec:analysis} we describe our data selections, the {\\em Fermi} LAT AGN and GRB observations during the first year of operation and analysis, and we discuss potential biases in the selection.  Our methodology and results are presented in Section~\\ref{sec:methods}.  We discuss implications of our results in Section~\\ref{sec:discussion}, and conclude in Section~\\ref{sec:conclusion}.  In the following, energies are in the observer frame except where noted otherwise. ",
        "conclusions": "\\label{sec:conclusion}  Using the high-energy 11-month photon data set collected by {\\em Fermi} from distant blazars, and two GRBs we have (i) placed upper limits on the opacity of the Universe to $\\g$ rays in the $\\sim$10--100~GeV range coming from various redshifts up to $z\\approx 4.3$; and (ii) ruled out an EBL intensity in the redshift range $\\sim$ 1 to 4.3 as great as that predicted by \\citet{Stecker06} in the ultraviolet range at more than $4\\sigma$ post trials in two independent sources (blazars). The overall rejection significance is found to be $>10 \\sigma$ post trials therefore making this result very robust. Our most constraining sources are blazars J1504+1029, J0808-0751 and J1016+0513 with $(z, \\langle E_{max} \\rangle)$ combinations of (1.84, 48.9~GeV), (1.84, 46.8 GeV) and (1.71, 43.3 GeV), respectively. Although a likelihood ratio analysis of the latter source indicates that the sensitivity of our analysis method is approaching the EBL flux level of the ``high UV model'' of Kneiske et al (2004), multi-trial effects markedly reduced the rejection significance. The two most constraining GRBs are GRB 090902B and GRB 080916C, both of which rule out the ``baseline'' EBL model of \\citet{Stecker06} in the UV energy range at more than 3$\\sigma$ level. The ``fast evolution'' model of \\citet{Stecker06} predicts higher opacities in the LAT energy range at all redshifts, and therefore is also ruled out. Together with the results from VHE observations (e.g., \\citet{aharonian07,Mazin07}) the models by \\citet{Stecker06} seem now disfavored in the UV and mid-IR energy range. We have also calculated model-independent optical depth upper-limits $\\tau_{\\g\\g, \\rm UL} (z, \\langle E_{max} \\rangle)$ at 95\\% CL in the redshift $z\\simeq 1-2.1$ and $E_{\\rm max}\\approx 28-74$GeV ranges.  As the high-energy photon data set collected by {\\em Fermi} grows in the future and more blazars and GRBs are detected at constraining energies, the $(E, z)$ phase space that constrains $\\tau_{\\g\\g}$ will become more populated. This will provide us with unique opportunities to constrain the opacity of the Universe to $\\g$-rays over a large energy and redshift range, and eventually help us further understand the evolution of the intensity of the extragalactic background light over cosmic time."
    },
    "1005/1005.1959_arXiv.txt": {
        "abstract": "{Using photometry of NGC 1097 from the {\\em Herschel} PACS (Photodetector Array Camera and Spectrometer) instrument, we study the resolved properties of thermal dust continuum emission from a circumnuclear starburst ring with a radius $\\sim 900$ pc.  These observations are the first to resolve the structure of a circumnuclear ring at wavelengths that probe the peak (i.e.  $\\lambda \\sim 100$ \\micron) of the dust spectral energy distribution.  The ring dominates the far-infrared (far-IR) emission from the galaxy---the high angular resolution of PACS allows us to isolate the ring's contribution and we find it is responsible for 75, 60 and 55\\% of the total flux of NGC 1097 at 70, 100 and 160 \\micron, respectively.  We compare the far-IR structure of the ring to what is seen at other wavelengths and identify a sequence of far-IR bright knots that correspond to those seen in radio and mid-IR images.  The mid- and far-IR band ratios in the ring vary by less than $\\pm 20$\\% azimuthally, indicating modest variation in the radiation field heating the dust on $\\sim 600$ pc scales.  We explore various explanations for the azimuthal uniformity in the far-IR colors of the ring including a lack of well-defined age gradients in the young stellar cluster population, a dominant contribution to the far-IR emission from dust heated by older ($> 10$ Myr) stars and/or a quick smoothing of local enhancements in dust temperature due to the short orbital period of the ring.  Finally, we improve previous limits on the far-IR flux from the inner $\\sim 600$ pc of NGC 1097 by an order of magnitude, providing a better estimate of the total bolometric emission arising from the active galactic nucleus and its associated central starburst.} ",
        "introduction": "The central regions of galaxies host some of the most intense star-formation that we can observe in the local Universe in circumnuclear starburst rings.  Starburst rings are believed to be the consequence of the pile-up of inflowing gas and dust, driven by a non-axisymmetric potential from a stellar bar, on orbits located near the Inner Lindblad Resonance of the bar \\citep{combes85,athanassoula92}.  The high surface densities that exist in the ring lead to high star-formation rates. Indeed starburst rings are one of few regions in non-interacting galaxies where the formation of ``super star clusters'' commonly occurs \\citep{maoz96}. The stars formed in the ring can be numerous enough to drive the structural evolution of the galaxy \\citep{norman96,kormendy04} and can be the dominant power source for the galaxy's infrared (IR) emission.  Star-formation in circumnuclear rings occurs under conditions not normally found in the disks of galaxies: in addition to their high gas surface densities, these regions have dynamical timescales that are comparable to the lifetimes of massive stars.  Understanding star formation in circumnuclear rings has been a long-standing problem \\citep{combes96}.  There are two main models: the ``popcorn'' model \\citep{elmegreen94}, where star-formation is driven by stochastic gravitational fragmentation along the ring, and the ``pearls on a string'' model, where gas flowing into the ring is compressed near the contact points (i.e. locations where the dust lanes intersect the ring) and then forms stars a short distance downstream \\citep[e.g.,][]{boker08}.  The ``pearls on a string'' model predicts a gradient in the ages of young stellar clusters as one moves away from the contact points.  This has been observed in a number of starburst rings \\citep[e.g., ][]{mazzuca08,boker08}.  Conversely, many well-studied rings show no evidence for an age gradient \\citep{maoz01}.  It is not obvious, however, that a single mode of star-formation must occur in all rings or even at all times in a given ring \\citep{vandeven09}.  KINGFISH (Key Insights into Nearby Galaxies: A Far-Infrared Survey with {\\em Herschel}, PI R. Kennicutt) is an Open-Time Key Program to study the interstellar medium (ISM) of nearby galaxies with far-IR/sub-mm photometry and spectroscopy.  Among the unique aspects of the KINGFISH science program is the ability to observe thermal dust emission at unprecedented spatial resolution ($\\sim$ 5.6, 6.8 and 11.3\\arcsec at 70, 100 and 160 \\micron) using PACS (Photodetector Array Camera and Spectrometer) imaging.  High spatial resolution is crucial for observing processes occurring in the central regions of galaxies.  These regions represent our best opportunity to study in detail the interplay between dynamics, star-formation and feedback that regulate the fueling of nuclear activity, be it a starburst or an active galactic nucleus (AGN).  Below we present PACS imaging of the galaxy NGC 1097, one of the first KINGFISH targets observed during the {\\em Herschel} Science Demonstration Program (SDP) (for PACS spectroscopy of NGC 1097 see Beir\\~{a}o et al. 2010 and for {\\em SPIRE} observations see Engelbracht et al. 2010).  The source NGC 1097 is a barred spiral galaxy located at a distance of 19.1 Mpc \\citep[][1\\arcsec$\\approx 92$ pc]{willick97}.  In its central kpc it hosts an intensely star-forming \\citep[$\\sim 5$ \\msun\\ yr$^{-1}$;][]{hummel87} ring with a radius of $\\sim 900$ pc.  The ring's rotation speed of $\\sim 300$ km s$^{-1}$ \\citep[corrected for inclination,][]{storchi-bergmann96}, corresponds to a rotation period of $\\sim 18$ Myr.  The galaxy's nucleus is classified as a LINER from optical emission line diagnostics \\citep{phillips84}, but is shown to be a Seyfert 1 by its double-peaked H$\\alpha$ profile \\citep{storchi-bergmann93}.  UV spectroscopy has revealed a few Myr old burst of star-formation in the central 9 pc of the galaxy \\citep{storchi-bergmann05}.  With the high spatial resolution of {\\em Herschel} PACS, we can resolve the starburst ring and inner 600 pc of NGC 1097 for the first time at wavelengths near the peak of the dust spectral energy distribution (SED). ",
        "conclusions": "We have presented {\\em Herschel} PACS observations from KINGFISH of the inner kpc of the barred spiral galaxy NGC 1097.  These are the first observations to resolve a starburst ring at wavelengths probing the peak of the dust SED.  We show a comparison of the ring in a variety of tracers and find similar bright knots in the mid- and far-IR and radio continuum.  These knots do not correspond to the same knots traced by CO.  We find modest variation azimuthally in the mid- and far-IR band ratios suggesting that either there is no azimuthal age gradient, as would be predicted by the ``pearls on a string'' mode of star-formation, that dust heating is dominated by an older stellar population and/or that the dust heating variations get quickly erased over the short ring orbital period ($\\sim 18$ Myr). Finally, we place an order-of-magnitude tighter constraint on the far-IR emission originating in the central $\\sim 600$ pc of the galaxy."
    },
    "1005/1005.2165.txt": {
        "abstract": "The apparent accelerating expansion of the Universe is forcing us to examine the foundational aspects of the standard model of cosmology -- in particular, the fact that dark energy is a direct consequence of the homogeneity assumption. We discuss the foundations of the assumption of spatial homogeneity, in the case when the Copernican Principle is adopted. We present results that show how (almost-) homogeneity follows from (almost-) isotropy of various observables. The analysis requires the fully nonlinear field equations -- i.e., it is not possible to use second- or higher-order perturbation theory, since one cannot assume a homogeneous and isotropic background. Then we consider what happens if the Copernican Principle is abandoned in our Hubble volume. The simplest models are inhomogeneous but spherically symmetric universes which do not require dark energy to fit the distance modulus.  Key problems in these models are to compute the CMB anisotropies and the features of large-scale structure. We review how to construct perturbation theory on a non-homogeneous cosmological background, and discuss the complexities that arise in using this to determine the growth of large-scale structure. ",
        "introduction": "The standard model of the Universe -- the LCDM ``concordance\" model -- is a perturbed FLRW model containing cold dark matter and dark energy in the form of a cosmological constant $\\Lambda$. This model is highly successful in being able, up to now, to fit all cosmological observations, with the same small set of parameters~\\cite{1}. However, there is as yet no satisfactory explanation for the value of $\\Lambda$, which appears in Einstein's field equations in precisely the same form as the total vacuum energy of quantum fields. This term is responsible for the late-time acceleration of the expansion of the Universe within the spatially homogeneous FLRW framework.  The unresolved nature of $\\Lambda$ and of alternative forms of dark energy throws into sharp focus the foundations of the standard model -- in particular, the spatial homogeneity assumption. We will address two questions in this paper:  \\begin{itemize} \\item[{\\bf (Q1)} ] What is the basis for the spatial homogeneity assumption? \\item[{\\bf (Q2)} ] If we drop this assumption, can we find a model that is consistent with all cosmological observations? \\end{itemize}  In summary, the current answer to (Q1) is that there are various results which motivate homogeneity (assuming the Copernican Principle); the most important one is that the observed (almost-)isotropy of the CMB provides strong evidence for (almost-)homogeneity of the Universe on large scales~\\cite{2}. But crucial open questions remain, which we identify and discuss. Standard cosmological perturbation theory is powerless here -- the argument depends on a fully nonlinear analysis, since one cannot assume an FLRW background.  The current answer to (Q2) is that the simplest inhomogeneous cosmological models -- isotropic LTB (Lema\\^\\i tre-Tolman-Bondi) models without any dark energy, and with our galaxy at the centre of a large void~-- are able to fit current supernova and some other data~\\cite{MT1,MT2,MHE,PascualSanchez:1999zr,celerier1,Tomita:2000jj,IKN,Mof1,Mof2,celerier2,VFW,CR,EM,BMN,bolejko,enqvist,YKN,AAG,Alnes:2006pf,alnes,ABNV,celerier3,Romano:2007zz,gbh1,gbh2,Enqvist,ZMS,CFL,BW,CCF,CFZ,CBKH,FLSC,mortsell,Mof3,Yoo:2010qy,Romano:2009ej,Romano:2009mr,Romano:2009qx}. In this framework, the late-time acceleration becomes a mis-interpretation of an effect due to nonlinear inhomogeneity and curvature. In order to test these models against the full range of cosmological observations, it is necessary to develop  perturbation theory on an LTB background. Although standard cosmological perturbation theory provides some qualitative guidance, perturbations on an inhomogeneous background are fundamentally different -- in particular, there is no longer a simple separation of scalar, vector and tensor modes at first order. We review recent developments on this~\\cite{CCF}, and discuss remaining open questions. ",
        "conclusions": "Without a theoretical understanding of the value of the cosmological constant the concordance model remains phenomenological no matter how strongly observations appear to support the model. The concordance model suffers an important coincidence problem: that dark energy starts to dominate at around the time our solar system forms. Until we demonstrate observationally that the Universe is homogeneous on large scales, we should consider inhomogeneous spacetimes even if they are philosophically uncomfortable, particularly in light of the fact that in their simplest incarnation they can explain away the dark energy problem through inhomogeneity, without apparently causing other problems. It is worth reflecting on how remarkable this is. Measurements of a nonzero $\\Lambda$ may in fact be thought of as an explicit consequence of the Copernican assumption.   It is important therefore to re-examine the basic assumptions of the standard concordance model -- both to test the Copernican Principle at the heart of modern cosmology, and to understand what freedom we have to develop new cosmological models that may be able to explain cosmic acceleration using curvature and inhomogeneity rather than a cosmological constant. In order to do this, we need non-perturbative methods and also new types of nonlinear perturbations where the nonlinearity arises from the inhomogeneous background. Moreover, the whole inflationary paradigm needs to be re-examined to see if inhomogeneities this large could arise naturally.   We have shown how the Copernican assumption when combined with the high isotropy of the CMB implies FLRW under fairly weak assumptions~-- mainly, that the CMB rest frame is geodesic. Provided that the dark matter is CDM that is comoving with the CMB, this geodesic property will hold. But in interacting dark energy-dark matter models, for example, the geodesic radiation frame is an additional assumption.  In order to test the Copernican Principle at the foundation of the FLRW models, observations which look inside our past lightcone to estimate how distant observers see the CMB~\\cite{gbh2,goodman,CS,Jia:2008ti} are necessary. In order to observationally prove FLRW, however, we also must demonstrate that galaxies follow geodesics, and that dark energy is comoving with the CMB.  If galaxies were non-geodesic, then this would leave a signature as a dipole in the Hubble law which grows linearly with distance (and so is distinguishable from a normal bulk flow, which is constant)~\\cite{CB} \\iffalse : in an inhomogeneous model the Hubble law becomes~\\cite{KS} \\be H_0={\\frac{1}{3}\\Theta}+ {A_\\a e^\\a}+ {\\sigma_{\\a\\nu}e^\\a e^\\nu},~~ e_\\a u^\\a=0\\,,~e_\\a e^\\a=1\\,, \\label{hubble from series} \\ee \\fi which is derived from the term which is linear in redshift in the area distance-redshift relation~-- see Eq.~(\\ref{hubble from series}). %(Note that the number of $e$'s tell us the spherical harmonic $\\ell$: e.g., the shear term is a quadrupole.) Possible dark energy flux would leave a similar signal in the quadratic term in the area distance-redshift relation. The generalized deceleration parameter, in a fully inhomogeneous spacetime is, following~\\cite{Cthesis}, \\ba %\\fl %\\begin{array}{rl} H_0^2\\, q_0&=&\\;16\\rho+\\;12p-\\;13\\Lambda -\\;23\\D_\\a\\udot^\\a +\\;{12}{5}\\sigma_{\\a\\nu} \\sigma^{\\a\\nu}-\\;23\\omega_\\a \\omega^\\a \\nonumber\\\\&& +e^\\a\\Bigg[-\\;23\\Theta \\udot_\\a -\\;{3}{5}\\sdel_\\a\\Theta-\\dot A_\\a+\\udot^\\nu\\sigma_{\\a\\nu}+\\;25q_\\a \\nonumber\\\\ &&%\\left.\\left. ~~~~~~~~~+\\;35\\varepsilon_\\a^{~ \\nu\\gamma}\\left(\\sdel_\\nu\\omega_\\gamma +2\\udot_\\nu\\omega_\\gamma\\right)\\Bigg]%\\right. \\nonumber\\\\ && \\left.+e^\\a e^\\nu\\left[-2\\sdel_{\\<\\a}\\udot_{\\nu\\>}+\\udot_{\\<\\a}\\udot_{\\nu\\>} +E_{\\<\\a\\nu\\>}-\\;12\\pi_{\\<\\a\\nu\\>}+ \\;97\\sigma^{\\vphantom{c}}_{\\gamma\\<\\a} \\sigma_{\\nu\\>}^{\\phantom{\\nu\\>}\\gamma} \\right.\\right. \\nonumber\\\\ &&\\left.\\left. +\\omega_{\\<\\a}\\omega_{\\nu\\>}- \\;23\\sigma^{\\vphantom{c}}_{\\gamma\\<\\a} \\omega_{\\nu\\>}^{\\phantom{\\nu\\>}\\gamma} \\right]\\right. + e^\\a e^\\nu e^\\gamma\\left[-5\\udot_{\\<\\a}\\sigma_{\\nu\\gamma\\>}-\\sdel_{\\<\\a} \\sigma_{\\nu\\gamma\\>}\\right]\\nonumber\\\\ && -3e^\\a e^\\nu e^\\gamma e^\\delta\\sigma_{\\<\\a\\nu}\\sigma_{\\gamma\\delta\\>} .\\label{q_0-observable} %\\end{array} \\ea If it were confirmed observationally that the CMB frame is geodesic, and that the CMB is almost isotropic around distant observers, then this $H_0^2q_0$ should be an isotropic observable~-- a potentially useful consistency check.  Although it seems obvious that CDM and baryons are just geodesic dust, we have to be careful because of issues highlighted from the averaging problem. While a distribution of non-relativistic CDM particles obey a dust equation of state at all times, more and more of the CDM particles are locked up as galaxies and clusters form, and yet those virialized clusters then become the `dust particles', within the standard paradigm. Are these geodesic? It is not yet clear how to treat this situation properly, and it is conceivable that there is a backreaction effect which changes the effective equation of state of matter at late times~\\cite{BR} (amongst other effects). Depending on precisely what is calculated, this backreaction can be divergent in the concordance model~\\cite{CAL}, and remains a significant problem to calculate non-linearly (although see~\\cite{Baumann:2010tm} who find no such divergences using a different method of averaging).      Off-lightcone observations will be able to place limits on the spatial gradients of the CMB multipoles as well as the multipoles themselves, which we have shown is also a requirement for deducing almost-FLRW geometry from almost-isotropic CMB (the almost-EGS result). The important exact ETM result states that if the first three multipoles are zero then all the rest must be too. It seems reasonable that this would extend to an almost-ETM result if the first 3 multipoles are small. In this case observations which look inside our past lightcone might only need to confirm that the first three multipoles around distant observers are small in order to observationally confirm FLRW. These issues deserve further investigation, and are important because it might be possible to construct a genuinely inhomogeneous model which solves the dark energy problem, but which also satisfies the Copernican principle, maintaining the isotropy of the CMB for all or `most' observers (as claimed in~\\cite{Wiltshire:2007zj}, for example).    %All results here rely on general relativity being the correct %description of gravity. In modified gravity, it is reasonably %straightforward to generalize to those cases too by treating the %extra gravitational degrees of freedom as effective matter terms %(e.g., in scalar-tensor gravity the EGS theorem holds~\\cite{CCO}). %In this case, effective energy flux terms and anisotropic pressure %terms would need to be observed to be small from local observations %as well as the CMB.   Dropping the Copernican Principle, we have the key exact result that isotropy of matter observations down our past lightcone implies an LTB geometry. However, it is an open question whether almost-isotropy of matter observations on the past lightcone implies an almost-spherically symmetric LTB model. We have discussed a complete analysis of structure growth at first order, which is the next step in developing LTB models, and shown how this is much more involved than the linear perturbations of FLRW. However, we argue that it is worth the effort. All tests of the Copernican Principle, including those which can be done by examining the background dynamics down our past lightcone~\\cite{CBL,UCE}, are difficult and sensitive. Only by developing a family of inhomogeneous spacetimes to the same level of sophistication as the standard concordance model, and directly comparing them side by side, will we really be able to understand whether $\\Lambda$ is real, or actually a consequence of our homogeneity assumption.                      ~\\\\{\\bf Acknowledgments:} We thank Timothy Clifton, Ruth Durrer,  George Ellis, Pedro Ferreira, David Matravers, John Moffat, Syksy R\\\"as\\\"anen and Marco Regis for comments and/or discussions. CC is supported by the National Research Foundation (NRF, South Africa) and RM is supported by the UK Science \\& Technology Facilities Council. This work was partially supported by the South African National Institute for Theoretical Physics and by a Royal Society (UK)/ NRF exchange grant between the Universities of Portsmouth and Cape Town.  \\newpage \\appendix"
    },
    "1005/1005.0114_arXiv.txt": {
        "abstract": "We present a new code, CASTRO, that solves the multicomponent compressible hydrodynamic equations for astrophysical flows including self-gravity, nuclear reactions and radiation. CASTRO uses an Eulerian grid and incorporates adaptive mesh refinement (AMR). Our approach to AMR uses a nested hierarchy of logically-rectangular grids with simultaneous refinement in both space and time. The radiation component of CASTRO will be described in detail in the next paper, Part II, of this series. ",
        "introduction": "In this paper, Part I of a two-part series, we present a new code, CASTRO, that solves the multicomponent compressible hydrodynamic equations with a general equation of state for astrophysical flows. Additional physics include self-gravity, nuclear reactions, and radiation. CASTRO uses an Eulerian grid and incorporates adaptive mesh refinement (AMR).  Our approach to AMR uses a nested hierarchy of logically-rectangular grids with simultaneous refinement of the grids in both space and time.  Spherical (in 1D), cylindrical (in 1D or 2D), and Cartesian (in 1D, 2D or 3D) coordinate systems are supported. The radiation component of CASTRO will be described in detail in the next paper, Part II, of this series.  There are a number of other adaptive mesh codes for compressible astrophysical flows, most notably, ENZO \\citep{ENZO}, FLASH \\citep{flash}, and RAGE \\citep{RAGE}. CASTRO differs from these codes in several ways.   CASTRO uses an unsplit version of the piecewise parabolic method, PPM, with new limiters that avoid reducing the accuracy of the scheme at smooth extrema; the other codes are based on operator-split hydrodynamics, though the most recent release of FLASH, version 3.2, includes an unsplit MUSCL-Hancock scheme. The different methodologies also vary in their approach to adaptive mesh refinement. RAGE uses a cell-by-cell refinement strategy while the other codes use patch-based refinement. FLASH uses equal size patches whereas ENZO and CASTRO allow arbitrary sized patches. ENZO and FLASH enforce a strict parent-child relationship between patches; i.e., each refined patch is fully contained within a single parent patch; CASTRO requires only that the union of fine patches be contained within the union of coarser patches with a suitable proper nesting. Additionally, FLASH and RAGE use a single time step across all levels while CASTRO and ENZO support subcycling in time.   All four codes include support for calculation of self-gravity.  It is worth noting that CASTRO uses the same grid structure as the low Mach number astrophysics code, MAESTRO (see, e.g., \\cite{multilevel}). This will enable us to map the results from a low Mach number simulation, such as that of the convective period and ignition of a Type Ia supernova, to the initial conditions for a compressible simulation such as that of the explosion itself, thus taking advantage of the accuracy and efficiency of each approach as appropriate. ",
        "conclusions": "We have described a new Eulerian adaptive mesh code, CASTRO, for solving the multicomponent compressible hydrodynamic equations with a general equation of state for astrophysical flows.  CASTRO differs from existing codes of its type in that it uses unsplit PPM for its hydrodynamic evolution, subcycling in time, and a nested hierarchy of logically-rectangular grids.  Additional physics includes self-gravitation, nuclear reactions, and radiation. Radiation will be described in detail in the next paper, Part II, of this series.  CASTRO is currently being used in simulations of Type Ia supernovae and core-collapse supernovae; examples of simulations done using CASTRO can be found in \\citet{Joggerstetal:2009,woosley-scidac2009}. Further details on the CASTRO algorithm can be found in the CASTRO User Guide \\citep{CASTROUserGuide}."
    },
    "1005/1005.0863_arXiv.txt": {
        "abstract": "We present a model for the rotational evolution of a young, solar mass star interacting with an accretion disk. The model incorporates a description of the angular momentum transfer between the star and disk due to a magnetic connection, and includes changes in the star's mass and radius and a decreasing accretion rate.  The model also includes, for the first time in a spin evolution model, the opening of the stellar magnetic field lines, as expected to arise from twisting via star-disk differential rotation.  In order to isolate the effect that this has on the star-disk interaction torques, we neglect the influence of torques that may arise from open field regions connected to the star or disk.  For a range of magnetic field strengths, accretion rates, and initial spin rates, we compute the stellar spin rates of pre-main-sequence stars as they evolve on the Hayashi track to an age of 3~Myr.  How much the field opening affects the spin depends on the strength of the coupling of the magnetic field to the disk.  For the relatively strong coupling (i.e., high magnetic Reynolds number) expected in real systems, all models predict spin periods of less than $\\sim3$ days, in the age range of 1--3~Myr. Furthermore, these systems typically do not reach an equilibrium spin rate within 3~Myr, so that the spin at any given time depends upon the choice of initial spin rate.  This corroborates earlier suggestions that, in order to explain the full range of observed rotation periods of approximately $1$--$10$ days, additional processes, such as the angular momentum loss from powerful stellar winds, are necessary. ",
        "introduction": "\\label{sec_intro}    The last three decades of observational surveys of pre-main-sequence stars have resulted in the measurement of rotation periods and/or rotational velocities ($v \\sin i$) of thousands of stars \\citep[see reviews by][]{herbstea07, scholz09}.  The analysis of these data have revealed a number of mysteries regarding the spin rates of young stars.  Perhaps the most prominent of these mysteries is that which was first noted by \\citet{vogelkuhi81}.  Specifically, a large fraction of nearly solar mass pre-main-sequence stars rotate much more slowly than expected.  The fact that stars are built up from material with high specific angular momentum, and that these young stars are still contracting, leads to the expectation of a rotation rate near the breakup velocity.  Solar mass stars with ages less than a few Myr have rotation periods typically in the range of 1--10 days but with a tail in the distribution extending to around 20 days \\citep[e.g.,][]{attridgeherbst92, choiherbst96, stassunea99, rebull01, rebull3ea04, coveyea05, scholz09}.  The statistics show that about half of these stars are indeed fairly rapid rotators and do seem to spin up as they approach the main sequence \\citep{bouvier3ea97, delarezapinzon04, scholzea07, irwinea08}.  However, roughly half of the stars younger than a few Myr rotate at about 10\\% or less of breakup speed \\citep[][]{rebull3ea04, herbstea07, scholz09}.  Thus, there is some mechanism operating that is capable of removing significant amounts of angular momentum during the pre-main-sequence phase.  Following the suggestion of \\citet{edwardsea93}, the modeling work of \\citet{bouvier3ea97} and \\citet{rebull3ea04} showed that the observed range of spins could be reproduced reasonably well, by assuming that the presence of an accretion disk somehow results in a constant stellar spin period of around one week.  There is some observational evidence that the population of stars with disks on average rotate more slowly than those without disks \\citep[e.g.,][]{edwardsea93, choiherbst96, stassunea99, herbstea00, littlefairea05, rebullea06, ciezabaliber07}.  Also, the idea that the angular spin rate is held at a constant by the presence of the disk is loosly based on physical models for the interaction between a star and surrounding disk.  There are generally two prominent theoretical ideas for how accreting stars can maintain a slow spin rate.  One idea is that the torques arising from the magentic connection between the star and disk can remove substantial angular momentum \\citep[e.g.][]{ghoshlamb78, camenzind90, konigl91, shuea94}.  When these torques are strong enough to enforce an equilibrium stellar spin rate, this idea is generally referred to as ``disk locking'' \\citep{choiherbst96}.  The other idea to explain slowly rotating accretors is that powerful stellar winds are primarily responsible for removing angular momentum from the star \\citep{hartmannmacgregor82, mestel84, hartmannstauffer89, toutpringle92, paatzcamenzind96, mattpudritz05l, mattpudritz08II, mattpudritz08III}.  Numerical, dynamical simulations of the star-disk interaction \\citep[such as the investigations of spin equilibrium by][]{romanovaea02, long3ea05} typically exhibit significant torques arising both from the magnetic star-disk connection and from winds.  The goal of the present paper is to further examine the possibility of disk locking as an explanation for the slow rotation of young stars. To this end, we develop a model for the time-evolution of stellar spin rates, under the influence of magnetic star-disk interaction torques. In this work, we neglect the influence of stellar winds, as is customary in the disk-locking models.  There are generally two types of disk locking models. One is the X-wind model \\citep[e.g.,][]{shuea94, ostrikershu95, mohantyshu08}. This model is developed under the assumption that the star-disk system on average exists in a spin equilibrium state, in which the accreting star feels no torque.  Thus, the X-wind model cannot be used to address whether or how the system reaches this equilibrium, nor how the system behaves when out of equilibrium (e.g., due to variability or due to the evolution of the star or disk).  All other disk locking models \\citep[e.g.,][]{camenzind90, konigl91, wang95, mattpudritz04} are of the second type, based on the general picture developed by \\citet{ghoshlamb78} for accreting neutron stars.  The Ghosh \\& Lamb model provides a method for calculating the torque on the star, due to the star-disk interaction, for any state of the system.  The star may spin up or down, depending on conditions, and there is a hypothetical disk-locked state, in which the net torque on the star is zero.  Thus, the Ghosh \\& Lamb model can be used to compute the evolution of accreting star spins and to assess how far from equilibrium a system may be.  \\citet[][hereafter CC93]{cameroncampbell93}, \\citet{yi94, yi95}, and \\citet[][hereafter AC96]{armitageclarke96} computed the evolution of pre-main-sequence star spins, adopting a Ghosh \\& Lamb-type model for torques on the star, and also including stellar contraction, a decrease in accretion rate with time, and a prescription for the evolution of the magnetic field.  For parameters that reasonably represent T Tauri stars, their models produced spin rates within the observed range.  These models also showed that, regardless of the initial spin of the star, the spin rates rapidly (in less than $\\sim 1$~Myr) approached an equilibrium value and stayed near equilibrium during the first few Myr of evolution.  That is, the initial spin conditions were ``erased,'' and the stellar spin was ``locked'' to a rate that depended only on the stellar mass and radius, magnetic field strength, and accretion rate.  These results generally supported the idea that disk locking explains the slow rotation of accreting T Tauri stars.  However, since that work, it has been recognized that there is a serious theoretical problem with a key assumption of the classical Ghosh \\& Lamb torque model.  The assumption is that the stellar magnetic field lines connect to a large region of the disk and become highly twisted in the azimuthal direction.  This twisting is what gives rise to the magnetic spin-down torque felt by the star.  Several authors \\citep[e.g.,][]{vanballegooijen94, lyndenbellboily94, lovelace3ea95, hayashi3ea96, millerstone97, goodson3ea97, fendtelstner00, mattea02, uzdensky3ea02, romanovaea02, kuker3ea03, bessolazea08} have now shown that, when the dipole field is twisted azimuthally past a threshold of approximately $45^\\circ$, the magnetic pressure associated with the azimuthal component pushes the field outward.  This leads to an inflation and opening of the field loops in a stellar rotation timescale, ultimately disconnecting the star from the disk.  \\citet{uzdensky3ea02} developed a method for computing which magnetic field lines would open.  The amount of field opening depends primarily on how strongly coupled the magnetic field is to the disk, since the coupling is responsible for the twisting, and the twisting is responsible for the field opening.  In order to assess how the field opening affects the torques, \\citet[][hereafter MP05]{mattpudritz05} modified the Ghosh \\& Lamb formulation to include the effects of field opening, following the method of \\citet{uzdensky3ea02}.  MP05 showed that, for the relatively strong coupling expected in these systems, the opening of the field significantly reduces the spin-down torque on the star, relative to the classical assumption of a closed field.  Thus, the equilibrium spin rate predicted by the disk locking picture is much faster, which calls into question whether disk locking can explain the existence of slowly rotating accretors.  However, since the MP05 model could only calculate the torque for a given set of parameters (at a given epoch), their conclusion is based on the calculation of the equilibrium spin rate.  MP05 were not able to describe how the spin evolves in time.  Therefore, there is a question as to whether the the stars actually evolve near equilibrium, and whether the evolution of the system, including changes in the accretion rate, stellar radius, and stellar spin rate may affect this conclusion.  Also, it is important to develop a physical model for understanding the evolution of stellar spin over a range of ages.  For these reasons, in this paper, we develop and utilize a stellar spin evolution model \\citep[similar to CC93; AC96;][]{yi94, yi95} that includes the updated torque theory of MP05.  In this first effort, we do not attempt to explain all phenomena related to young star spins. Rather, we focus on testing whether models can produce spin rates within the observed typical range of $\\sim1$--10 days.  This work extends the MP05 formulation to the time domain, during the first $3\\times 10^{6}$ yr (i.e., during the Hayashi phase).  Section \\ref{sec_model} contains a description of the model and details of our calculations.  Section \\ref{sec_results} contains the results of stellar spin calculations that include the effect of magnetic field opening.  A discussion and summary of the conclusions of this work are contained in section \\ref{sec_conclusion}.  As a test of the model, and for illustrative purposes, we include an Appendix containing the results of stellar spin calculations under the classical assumption of a completely closed field. ",
        "conclusions": "\\label{sec_conclusion}    The star-disk interaction model of MP05 includes the effects of the field opening expected to arise from star-disk differential rotation ($\\gamma_c=1$), as well as the effects of magnetic coupling, parameterized by $\\beta$.  Based essentially upon the idea that the disk likely exists in a state of high magnetic Reynolds number, MP05 argued that a value of $\\beta=0.01$ is appropriate for real systems. In the present work, we have extended the MP05 torque model into the time domain, in order to explore the consequences of magnetic star-disk coupling and connectedness on the evolution of stellar spins.  To this end, we developed a stellar spin evolution model (\\S \\ref{sec_model}).  The model follows an accreting, one solar mass star during the Hayashi phase (from $3\\times10^4$ yr to $3$ Myr) and considers the torques expected to arise from the magnetic star-disk interaction.  The primary goals of this work were to examine the role of disk locking and to determine if the star-disk interaction torques alone are sufficient to explain the observed range of spin rates.  In developing our torque and spin-evolution model, we have adopted many of the same assumptions as the disk-locking models in the literature, except that the MP05 torque formulation includes the effect of the opening of the magnetic field.  Thus, while there are inherant uncertainties associated with the adoption of many of these assumptions, our results serve best to highlight the effect of field opening {\\it relative} to previous results in the literature.  Given the expectation that $\\gamma_c=1$ and $\\beta=0.01$ best-represents the conditions in real systems, and for a range of accretion rates and field strengths appropriate for T Tauri stars, the two main conclusions of the present work are as follows:  1. The models presented in figures \\ref{fig_g1b500bet01} and \\ref{fig_g1b2000bet01} exhibit spin periods ranging from 3 days to less than 1/3 day, in the age range of 1--3~Myr (see \\S \\ref{sec_gam1}).  These are consistent only with the fastest rotators in the observed spin period distribution, which is significantly populated from approximately 1--10~days (see \\S \\ref{sec_intro}).  It is apparent that the torque arising from the magnetic connection between the star and disk is not sufficient to explain the relatively slowly spinning, young stars.  2. Furthermore, the stars in these models are generally not in spin equilibrium (with a net zero torque), and the spin rate at all times depends upon the initial spin rate (see \\S \\ref{sec_gam1}).  Only the models with the strongest field strength and highest accretion rate neared their equilibrium spin rate in $\\sim1$~Myr (see figure \\ref{fig_g1b2000bet01}).  However, this equilibrium had a relatively short spin period of less than a day.  It is apparent that disk locking does not play a strong role, except possibly for the most rapid rotators.  Note that the term ``disk locking'' generally refers to the idea that the angular spin rate of the star is nearly equal to that of the inner edge of the disk (i.e., $R_t \\approx R_{\\rm co}$).  In our magnetic models, this condition was always true.  In fact, the truncation radius was nearest to the corotation radius in the models that were furthest from spin equilibrium (e.g., compare panel (d) in figures \\ref{fig_g1b500bet01} and \\ref{fig_ginfb500bet1}).  Thus, these stars do have an angular spin rate that is very close to that of the disk inner edge, although it is not appropriate to think of these systems as being ``locked'' to any particular spin rate.  Instead, this is a consequence of how the disk is truncated.  When the magnetic coupling in the disk is strong (small $\\beta$), the disk will generally be truncated close to the corotation radius (in State 2; see \\S \\ref{sec_rt}).  But, the condition that $R_t \\approx R_{\\rm co}$ does not necessarily mean that the star experiences a net zero torque.  In order to explain the existence of slowly rotating young stars, it is necessary to consider other possibilities.  Since the magnetic torques depend strongly on the magnetic field strength, it is tempting to suggest that a stronger magnetic field may solve the problem. Using our model, we find that (for $\\beta=0.01$) a dipole field strength of $B_*=10^4$~G is required to maintain spin periods consistent with the slow rotators. However, no T Tauri star has been found to have surface field strengths greater than 3~kG \\citep{johnskrull07}, and the global (dipole) fields are generally even weaker \\citep[e.g.,][]{safier98, bouvierea07, donatiea07, donatiea08}.  Given the relatively small number of stars for which there are magnetic field measurements, the present situation could be improved by more and improved observations of the global field strength and geometry.  Since the models presented in section \\ref{sec_gam1} with weak magnetic coupling (figure \\ref{fig_g1b2000bet1}) exhibit slow spins, it is also appropriate to consider the possibility that real systems have weak coupling.  As discussed above, the conditions that are expected to best represent young star-disk systems from ``first principles'' suggest values of $\\beta=0.01$.  However, $\\beta$ is a highly uncertain parameter because the physics of magnetic coupling is not well understood, nor are the conditions that influence the coupling (e.g., the ionization fraction of disk gas, turbulence levels in the disk, or magnetic reconnection rates).  In order for magnetic coupling to explain the slow rotators, the coupling strength must not be very different (neither larger nor smaller) than $\\beta=1$, since MP05 showed that the maximum spin-down torque occurs for a $\\beta$ of unity.  The condition that $\\beta \\approx 1$ seems to require some fine-tuning, which makes this possiblity even more difficult to justify at present.  Therefore, it appears necessary to consider effects that are alternative or additional to the magnetic torques arising from the star-disk connection.  Real systems seem to be variable on all timescales \\citep[e.g.,][]{hartmann97}, and we have neglected variability here.  It may be possible, for example, that short timescale (e.g., $\\la 10^4$~yr), large-magnitude variations of the accretion rate are important for the stellar mass and angular momentum evolution \\citep[e.g.,][]{popham96}.  Finally, it is necessary to consider the angular momentum loss from pre-main-sequence stellar winds (see \\S \\ref{sec_intro}).  The idea that stellar winds may be important during the accretion phase is well-supported \\citep[e.g.,][]{decampli81, hartmann3ea82, kwantademaru88, hartmannea90, fendt3ea95, fendtcamenzind96, safier98, bogovalovtsinganos01, sauty3ea02, edwardsea03, edwardsea06, dupreeea05, meliani3ea06, kurosawa3ea06, kwan3ea07, fendt09}. Stellar winds may be powered by the accretion process itself \\citep{toutpringle92, cranmer08} and be the key driver of angular momentum loss \\citep{hartmannmacgregor82, mestel84, hartmannstauffer89, paatzcamenzind96, mattpudritz05l, mattpudritz08II, mattpudritz08III}.  Some recent magnetohydrodynamic (MHD) simulations of rapidly rotating stars \\citep[e.g.,][]{romanovaea05, romanovaea09} exhibit a type of ``propeller'' regime in which intermittent accretion occurs on typical timescales of several orbits of the inner disk, while most of the would-be accreting material is instead launched into a wind.  For the calculations in the present work, the torque represents that which is averaged over a calculation timestep.  Our timesteps range from approximately 400 years (at the earliest times) to $\\sim 10^5$ yrs. While some of the simulations in the work cited above have been run for thousands of dynamical times, this is still orders of magnitude shorter than the typical timestep in the present work.  It is not clear whether the torque formulation we adopted accurately describes the time-average behavior of the accretion and closed field region in the propeller regime simulations.  In addition, while these MHD simulations reliably demonstrate many of the complexities and sometimes episodic nature of the magnetic star-disk interaction, when considering the long-term evolution of the star and disk, it is not yet clear whether the duration of the episodic (e.g., propeller) regime is significant for the overall spin evolution of young stars. However, the MHD simulations such as those cited above often exhibit significant outflows and torques arising from open magnetic field regions.  In this respect, our result that the torques arising only from closed field regions are not sufficient to significantly spin-down the star, agrees with the simulations.  It is clear that the opening of field lines is an important effect, since including this in the models indicates considerably different physics at work than the closed field models, when comparing to the same observational data.  In future work, we will extend the present analysis to include the angular momentum loss from accretion-powered stellar winds."
    },
    "1005/1005.5441_arXiv.txt": {
        "abstract": "We analyze the most salient cosmological features of axions in extensions of the Standard Model with a gauged anomalous extra $U(1)$ symmetry. The model is built by imposing the constraint of gauge invariance in the anomalous effective action, which is extended with Wess-Zumino counterterms. These generate axion-like interactions of the axions to the gauge fields and a gauged shift symmetry. The scalar sector is assumed to acquire a non-perturbative potential after inflation, at the electroweak phase transition, which induces a mixing of the St\\\"uckelberg field of the model with the scalars of the electroweak sector, and at the QCD phase transition. We discuss the possible mechanisms of sequential misalignments which could affect the axions of these models, and generated, in this case, at both transitions. We compute the contribution of these particles to dark matter, quantifying their relic densities as a function of the St\\\"uckelberg mass. We also show that models with a single anomalous U(1) in general do not account for the dark energy, due to the presence of mixed $U(1)-SU(3)$ anomalies. ",
        "introduction": "Given its important role as a possible solution of the strong CP problem \\cite{Peccei:1977ur} as well as a candidate for the dark matter of the universe, the study of axions \\cite{Wilczek:1977pj,Weinberg:1977ma} \\cite{Dine:1981rt,Zhitnitsky:1980tq,Kim:1979if,Shifman:1979if} (see \\cite{Sikivie:2006ni} for an overview) has received momentum both at theoretical and experimental level along the years. The invisible axion owes its origin to a global $U(1)_{PQ}$  (Peccei-Quinn, PQ) symmetry which is spontaneously broken in the early universe and explicitly broken to a discrete $Z_N$ symmetry by instanton effects at the QCD phase transition \\cite{Sikivie:1982qv}. The breaking occurs at a temperature $T_{PQ}$ below which the symmetry is nonlinearly realized. Strings and domain walls relics, which are typical of axion models and are a problem in ordinary PQ cosmology, can be avoided by introducing inflation to account for their dilution, or by embedding the model into more general constructions based on theories of Grand Unification \\cite{Lazarides:1982tw}.  The almost massless nature of the axion and its suppressed coupling to the fields of the Standard Model are consequences of the fact that this field is associated with the phase of a global anomalous symmetry. Both properties are related to the same scale, the axion decay constant $f_a\\sim 10^{10}-10^{12}$ GeV.  The implications of the PQ axion in cosmology, both in supersymmetric and in non supersymmetric models, have been explored to a finer level of detail. For instance, the axion plays an important role in determining the structure of the primordial perturbations \\cite{Dimopoulos:2003ii,Dimopoulos:2003ss,Lazarides:2005ek}, where it can act as a curvaton.  The gauging of an anomalous symmetry has some important effects on the properties of this pseudoscalar, first among all the appearance of independent mass and couplings to the gauge fields. This scenario allows a wider region of parameter space where to look for these particles. For this reason, axion-like fields, which are at the center of several investigations, are unlikely to find any significant and fundamental formulation without an underlying anomalous $U(1)$ gauge symmetry, as emphasized in previous works~\\cite{Coriano:2006xh,Coriano:2007xg}, \\cite{Coriano:2009zh}.  So far only two complete models have been put forward for a consistent analysis of these types of particles, the MLSOM \\cite{Coriano':2005js} and the USSM-A \\cite{Coriano:2008xa}. The first of them is at the basis of the elaborations that we are going to provide in this work. Here we will be focusing on the phenomenological analysis of a scenario which is a direct consequence of the model introduced in \\cite{Coriano':2005js}, while more details on the supersymmetric construction will be discussed in a separate work.  Although the natural framework that motivates these constructions is open string theory \\cite{Angelantonj:2002ct}, the effective actions describing these types of particles can be consistently defined at lower energy just by the inclusion of the relevant dimension-5 Wess-Zumino (Peccei-Quinn) interactions. These are necessary in order to guarantee the gauge invariance of the effective action and can be interpreted as counterterms. In fact, they balance the anomalous variation of the 1-loop effective action induced by the extra $U(1)$ symmetry, restoring the gauge symmetry.  The gauging of an anomalous symmetry is the essential element in the construction of these effective actions and can be justified within intersecting brane models. The gauging is a variant of the standard Peccei-Quinn construction and is characterized by a new scale $M$, which is the St\\\"uckelberg mass. We recall that St\\\"uckelberg extensions of the Standard Model with a non-anomalous $U(1)$ have been analyzed in several recent works \\cite{Feldman:2007wj,Feldman:2006wb,Feldman:2010wy}.  We are going to provide a physical perspective on the possible phenomenological implications of the anomalous case. In particular, we will try to connect the St\\\"uckelberg fields, which are in the spectrum of these models, to the physical axion which may appear as an extremely weakly interacting particle of a certain relic density in our current universe. Our assumption, in the identification of the physical axion, is that the original St\\\"uckelberg fields will mix at the electroweak phase transition with the Higgs sector. As a result, an almost massless state will emerge after the electroweak phase transition.  One of the key mechanisms that we will try to adapt and extend from the PQ case is that of {\\em vacuum misalignment}. This phenomenon occurs whenever a quasi Nambu-Goldstone mode - generated by the breaking of a certain symmetry - acquires non-perturbatively a small potential, lifting one flat direction from the vacuum degeneracy. For axions characterized both by an $SU(2)$ and an $SU(3)$ charge the mechanism of vacuum misalignment becomes sequential, as we are going to show.  From a more general perspective, we will also try to characterize the possible role of these types of particles as quintessence axions. These appear in models where the axions remain decoupled from the gluonic sector and their mass is purely of electroweak origin (see for instance \\cite{Nomura:2000yk}). In this case one tries to exploit the Nambu-Goldstone nature of these particles.  The main idea behind this proposal is that a phase transition around the electroweak scale can generate a small curvature in the potential, capable of giving a tiny mass to this particle, smaller than the Hubble parameter at current time ($H_0$). For this to be possible, as we are going to show, one has to search for solutions of the anomaly equations for an anomalous $U(1)$ which has a vanishing mixed anomaly with the $SU(3)$ color group. In this case the only source of mass for these axions would come from the electroweak and  not from the QCD phase transition, and as such could be extremely small.  In the general models that we analyze, the anomaly equations do not allow for such a solution, although this would not exclude the possibility of finding others, in the presence of more complicated gauge structures, for instance in models with several U(1)'s. We will not address this specific point any further, leaving it as an option for future studies. Instead, we will concentrate on the general features of an axion-like field coming from a single anomalous $U(1)$ symmetry, characterized by the presence of mixed anomalies both with the $SU(2)$ and $SU(3)$ sectors. The phenomenological details of the model are rather intricate, and have been worked out before. For this reason we have summarized in the next section some of their salient features, which turn out to be necessary in order to proceed with a realistic estimate of the relic densities. This is the specific goal of our work. ",
        "conclusions": "We have discussed the most salient cosmological features of models containing gauged axions, obtained from the gauging of an anomalous symmetry.  The gauging allows to define a consistent theory for axion-like particles, which generalize many of the properties of PQ axions. They have appeared for the first time in the study of intersecting branes, but their features are quite generic. They are constructed as effective theories containing minimal gauge interactions which restore gauge invariance of the effective action in the presence of an anomalous $U(1)$ symmetry, and no further requirements. Differently from the PQ case, here there is no concept of an original PQ symmetry, broken at a very large scale, with the axion taking the role of a Goldstone mode that acquires a mass at the QCD phase transition. Rather, the physical axion emerges directly at the electroweak phase transition, when Higgs-axion mixing occurs. Being charged under $SU(3)$ and $SU(2)$, we have a sequential misalignment of this field, and we have quantified its relic density as a function of the St\\\"uckelberg mass. We have shown that only very large values of the St\\\"uckelberg mass cause a significant contribution of this type of axions to the current dark matter content of the Universe, which otherwise remains negligible. The absence of an original PQ-like potential has some implications at cosmological level, such as the absence of isocurvature perturbations, since the St\\\"uckelberg is not a physical mode before the electroweak phase transition, in particular at the time of inflation. This feature is due to the presence of a local gauge symmetry, realized in the St\\\"uckelberg form, which allows to absorb $b$  into the longitudinal component of the anomalous gauge boson.  Our analysis represents, more generally, a description of the fate of the St\\\"uckelberg field in cosmology, from the defining St\\\"uckelberg phase of the theory at  a large scale (defined by the value of the St\\\"uckelberg mass), down to the electroweak and QCD phase transitions, when this field develops a physical component.  Our analysis could be extended in several directions, for instance with the inclusion of the modifications induced on the computation of the relics due to the presence of non-homogeneities in $\\chi$ beyond the QCD horizon, a feature which is also present in PQ models when the PQ scale lays below the scale of inflation. However, even at this level of refinement, the only significant scale in the determination of the relic densities remains the value of the St\\\"uckelberg mass.  Small values of this mass parameter in the TeV range leave the contribution of these particles to the relic densities of dark matter negligible, and sizeable for $M$ around an intermediate scale of $10^7$ GeV. In this case all the constraints coming from the neutral current sector are satisfied, being the extra $Z^\\prime$ of the theory completely decoupled from the low energy spectrum of the Standard Model. The appearance of this intermediate scale is a novel feature of this type of axions which could be used to set limits on their parameter space.  \\vspace{1cm}  \\centerline{\\bf Acknowledgements} We thank Pierre Sikivie and Nikos Irges for discussions. C.C. thanks the Physics Department at Thessaloniki for hospitality. This work is supported in part  by the European Union through the Marie Curie Research and Training Network UniverseNet (MRTN-CT-2006-035863). \\begin{appendix}"
    },
    "1005/1005.3784_arXiv.txt": {
        "abstract": "{We present a preliminary analysis of the small-scale structure found in new 70-520 \\um\\ continuum maps of the Rosette molecular cloud (RMC), obtained with the SPIRE and PACS instruments of the {\\it Herschel} Space Observatory. We find 473 clumps within the RMC using a new structure identification algorithm, with sizes up to $\\sim$1.0 pc in diameter.  A comparison with recent {\\it Spitzer} maps reveals that 371 clumps are ``starless\" (without an associated young stellar object), while 102 are ``protostellar.\" Using the respective values of dust temperature, we determine the clumps have masses ($M_{C}$) over the range -0.75 $\\leq$ log~($M_{C}$/\\msun) $\\leq$ 2.50. Linear fits to the high-mass tails of the resulting clump mass spectra (CMS) have slopes that are consistent with those found for high-mass clumps identified in CO emission by other groups.} ",
        "introduction": "\\label{intro}  Stars form within molecular clouds, after some fraction of cloud material is first condensed into smaller-scale structures.  How this process unfolds is not well understood, however.  For example, what conditions within clouds drive the formation of small-scale structures that in turn produce stars of any given mass, e.g., high-mass stars.  Detailed observations of the small-scale structure within molecular clouds should provide valuable insight into how stars of various masses form out of dense material.  The need for sensitive probes of small-scale structure in molecular clouds is the impetus behind the ``{\\it Herschel}\\/ OB Young Stellar objects\" (HOBYS) Key Project (see Motte et al. 2010).  The HOBYS team is currently using the ESA {\\it Herschel} Space Observatory (Pilbratt et al. 2010) to obtain wide-field maps of $\\sim$15 molecular clouds which are forming high-mass stars within 3 kpc of the Sun.  These clouds are being observed with both the {\\it Herschel} SPIRE (Griffin et al.\\/ 2010) and PACS (Poglitsch et al.\\/ 2010) instruments, ultimately to obtain wide maps of these clouds at 70-520 \\um\\ at diffraction-limited resolutions of $\\sim$18\\as\\ $\\times$ ($\\lambda$/250 \\um). Such data sample thermal emission from cold dust grains mixed with molecular gas within the cloud.  Since continuum emission at submillimetre wavelengths is typically optically thin, detections of continuum emission from dust at high resolution can trace the organization of mass in clouds on small scales.  By sensitively sampling emission across the peaks of the spectral energy distributions (SEDs) of this dust, SPIRE and PACS data can constrain both the column density and temperature of the dust and provide unparalleled censuses of the small-scale structure in molecular clouds.  In this paper, we describe the small-scale structure detected with SPIRE and PACS in the Rosette molecular cloud (RMC), the first of the HOBYS target list that was observed as part of the early {\\it Herschel} ``science demonstration phase\" (SDP) campaign.  The RMC is a rich location of star formation at a distance of 1.6 kpc (Johnson 1962; P\\'erez et al. 1987; Park \\& Sung 2002). The cloud is southeast of NGC 2244, the Rosette Nebula, and the expanding HII region associated with NGC 2244 has begun to interact with it.  The structure of material within the RMC has been previously examined by Williams, Blitz \\& Stark (1995), Schneider et al.\\/ (1998) and Dent et al.\\/ (2009) using various CO observations.  Alternatively, the populations of young stellar objects (YSOs) within the RMC have been studied by Phelps \\& Lada (1998), Li \\& Smith (2005), Poulton et al.\\/ (2008), and Rom\\'an-Z\\'u\\~niga et al.\\/ (2008) using infrared observations.  Notably, these studies identified 10 embedded clusters within the RMC, i.e., PL01-07 and REFL08-10.  A recent review of observations of the RMC has been made by Rom\\'an-Z\\'u\\~niga \\& Lada (2008; RZL).  The SPIRE and PACS observations described in this paper provide a preliminary look into how {\\it Herschel}\\/ data trace small-scale structures in the RMC, and the relationships between these structures and YSOs.  Analyses of other aspects within the SDP data of the RMC can be found in the accompanying papers by Motte et al.\\/ 2010, Henneman et al.\\/ 2010, and Schneider et al.\\/ 2010.) ",
        "conclusions": "\\label{summary}  We have conducted a preliminary analysis of small-scale emission within the Rosette molecular cloud, as observed by the {\\it Herschel} Space Observatory. From SPIRE and smoothed PACS images, we identified 473 clumps using {\\it getsources}, a new algorithm that identifies structure over multiple wavelengths and scales.  Using {\\it Spitzer}\\/ data, we classified these into 371 starless and 102 protostellar clumps.  The clumps have dust temperatures of 10-30 K, with starless clumps surprisingly warmer on average than the protostellar clumps, owing to the proximity of many to NGC 2244.  Masses were determined for each clump, revealing a range of -0.75 $\\leq$ log~($M_{C}$/\\msun) $\\leq$ 2.50.  Linear least-squares fits to the high-mass tails of the CMS reveal slopes that are consistent with slopes fit to mass spectra of clumps identified through CO in the RMC.  Such slopes are shallower than those seen in the high-mass tails of the mass spectra of small-scale continuum structures (``cores\") in closer clouds, possibly from crowding of such objects on scales smaller than probed here."
    },
    "1005/1005.3267_arXiv.txt": {
        "abstract": "{} {It was shown in previous works  the  existence of weakly chaotic orbits in the plutino population that diffuse very slowly. These orbits correspond to long-term plutino escapers and then represent the plutinos that are escaping from the resonance at present. In this paper we perform numerical simulations in order to explore the dynamical evolution of plutinos recently escaped from the resonance. } {The numerical simulations were divided in two parts. In the first one we evolved $20,000$ test particles in the resonance in order to detect and select the long-term escapers. In the second one, we numerically integrate the selected  escaped plutinos in order to study their dynamical post escaped behavior. } {Our main results include the characterization of the routes of escape of plutinos and  their evolution in the Centaur zone. We obtained a present rate of escape of plutinos between 1 and 10 every 10 years. The escaped plutinos have a mean lifetime in the Centaur zone of $ 108$ Myr and their contribution to the Centaur population would be a fraction of less than  $6 \\%$  of the total Centaur population. In this way, escaped plutinos would be a secondary source of Centaurs. } {} ",
        "introduction": "In the last few years the number of  observed transneptunian objects has enormously grown thanks to the progresses in the astronomical observations. This fact has allow to define, in a more rigorous way, the different dynamical classes previously identified in the first years of discoveries. The transneptunian region (TNR) can be structured into 4 dynamical classes  (Chiang et al. \\cite{Chiang07}). The Resonant Objects are those in mean motion resonance with Neptune, the Classical Objects are those non-resonant objects with semimajor axis $a$ greater than $\\sim$42 AU and low eccentricity orbits, the Scattered Disk Objects (SDOs) with perihelion distances $q > 30$ AU and large eccentricities, and the Centaurs Objects. This last group has perihelion distances inside the orbit of Neptune and they are transitory objects descendants of the other 3 classes, mainly from the SDOs, recently dislodge from the transneptunian zone by planetary perturbations (Levison \\& Duncan \\cite{Levison97},  Tiscareno \\& Malhotra \\cite{Tiscareno03}), Di Sisto \\& Brunini \\cite {Disisto07} ).  Centaurs are sometimes defined according to their aphelion distance or to their semimajor axis ($a$), as for example the nomenclature of Gladman et al  \\cite{Gladman08} that uses $a <a_N$ (where $a_N$ is Neptune's semimajor axis),  being then objects entirely in the giant planet zone. Nevertheless,  it is generally accepted that they are objects which enter the planetary region from the TNR, evolve in the giant planetary zone and a fraction of them enter the zone interior to the Jupiter's orbit becoming Jupiter Family Comets (JFCs).   The resonant transneptunian population most densely populated  are the plutinos, that are trapped into the 2:3 mean motion resonance with Neptune, being Pluto its most representative member. Some of the plutinos cross Neptune's orbit and hence  one might think  they would be sensitive to strong perturbations during close encounters with that planet. However conjunctions occur near plutinos's aphelions and then close encounters do not occur. This fact provides a stable configuration of the resonance, with the critical angle $\\sigma = 3 \\lambda - 2 \\lambda_N - \\varpi$ librating around $180^{\\circ}$, where $\\lambda$ and $\\varpi$ are the mean longitude and the longitude of perihelion of the plutino, and $\\lambda_N$ is the mean longitude of Neptune.  Duncan et al. (\\cite{Duncan95}) analyze  the  dynamical structure of the transneptunian region through numerical simulations.  They numerically integrate thousands of  test transneptunian bodies in order to  study their dynamical behavior and to determine which regions may be potential sources of the Jupiter family comets that we see today. With regards to  the Plutinos, Duncan et al. (\\cite{Duncan95}) showed that the boundaries of this long-lived mean motion resonance have a time scale for instability of the order of the  age of the Solar System, and so the long-term erosion of the particles is rather gradual and therefore must be continuing at the current epoch. The existence of these currently unstable orbits  may be related to the origin of the observed Jupiter family comets.  Morbidelli (\\cite{Morby97}) studied the dynamical structure of the 2:3 resonance with Neptune in order to analyze possible diffusive phenomena and their relation to the existence of long-term escape trajectories. He performed numerical simulations integrating the evolution of 150 test particles initially covering the 2:3 resonance with eccentricities up to $e = 0.3$ and inclinations less than $i = 5 ^ {\\circ }$, for 4 byr. The  author  found regular orbits that never escaped from the resonance present only at moderate eccentricity and small  amplitude of libration,  and  a  strong   chaotic  region  at  large  amplitude  of libration, which is quickly depleted and seems to be generated for the interaction  between the  $\\nu_{18}$ and  kozai resonances. Moreover, Morbidelli (\\cite{Morby97}) showed the existence of weakly chaotic orbits that diffuse very slowly and finally dive into the strong chaotic region. They are then long-term escapers and then plutinos recently escaped from the resonance. The source zone  of those particles   produce Neptune- encountering bodies at the current epoch of the Solar System and should be an active source of Centaurs and comets at present. He found that 10 $\\%$ of the bodies initially in that  weakly chaotic zone are delivered to Neptune in the last $10^9$ years.  The first  attempt at including  the gravitational influence  of Pluto into numerical models of the  dynamics of plutinos was developed by Yu \\& Tremaine (\\cite{Yu99}). These  authors suggested that the effect of Pluto,  ejecting objects  from the  2:3 resonance  to Neptune-crossing orbits, may contribute to or  even dominate the flux of JFCs. At the  same  time, Nesvorn\\'y  et al. (\\cite{Nesvorny00}) analyzed the effect of Pluto on the 2:3 resonant orbits. They found that  Pluto produces a large excitation of the libration amplitudes  in the 2:3 resonance. However they estimated that the flux rate that contributes to the flux of short period comets is about 1 $\\%$ of the 2:3 resonant population per $10 ^ {8 }$ years  which is about the same value as the flux obtained by Morbidelli  (\\cite{Morby97}) without Pluto.  In the early 2000s,  Melita \\& Brunini (\\cite{Melita00}) developed a comparative  study of  mean motion  resonances in  the transneptunian region. These  authors used the  frequency-map-analysis method (Laskar \\cite{Laskar93}) in  order to describe the dynamical  structure of the 2:3,  3:5 and  1:2  Neptune resonances,  at  39.5, 42.3  and 47.7  AU, respectively.  Particularly,  Melita  \\&  Brunini (\\cite{Melita00}) showed  that the  2:3 resonance  presents  a very  robust stable  zone primarily at  low inclinations, where  a great number of  the observed plutinos are distributed. They suggested that the existence of  plutinos in  very unstable  regions can  be explained  by physical collisions  or  gravitational   encounters  with  other  plutinos.  de El\\'{\\i}a et al. (\\cite{deelia08}) performed a collisional evolution of plutinos and obtained a plutino removal by their collisional evolution of 2 plutinos with $R>1$ km every 10000 years or a flux rate of escape of 0.5 $\\%$ of plutinos in $10^{10}$ years.  Very recently, Tiscareno \\& Malhotra (\\cite{Tiscareno09}) carried out 1-Gyr numerical integrations to study the characteristics of the 2:3 and 1:2 mean motion resonances with Neptune. Their main results include maps of resonance stability for a whole range of eccentricities and inclinations. They made integrations with and without Pluto, and concluded that it has only modest effects on the Plutino population. They calculated the fraction of remaining plutinos in the resonance as a function of time and extrapolated this fraction after 4 Gyr and also evaluated the fate of escaped particles.    From those previous works, plutino removal by ``dynamic'' is much greater than plutino removal by collisions. We refer to ``dynamical'',  those numerical simulations that take into account the gravitational forces to follow up the evolution of a particle and occasionally cause the removal. Then in this work, we perform ``dynamical'' numerical simulations in order to describe and characterize the routes of escape of plutinos and their contribution to the other minor bodies populations of the Solar System, specially to Centaurs. ",
        "conclusions": "We have performed two numerical simulations in order to first obtain particles representative to the plutinos that are escaping at present from the resonance and second in order to describe their dynamical post escape evolution. In the first simulation we integrate $20,000$ initial particles in the 2:3 resonance and find that $\\sim 88 \\%$ of the particles left out the integration and the rest remain in it. Considering a plutino population with radius grater than 1 km of $N_p \\sim 10^8 - 10^9$, we obtained a present rate of escape of plutinos between 1 and 10 every 10 years. From this integration we selected those particles that are representatives of the present escape plutinos and performed a second integration. From this last integration we obtained the dynamical evolution of plutinos once they escape from the resonance. From the 1183 initial particles, 1179  were removed from the integration and 4 remain in it. From the 1179  removed particles, 787 are ejected, 385 reached the Jupiter's zone and 4 collide with the planets.  We found that the great majority of escaped plutinos have encounters with Neptune, and this planet governs their dynamical evolution. When a plutino escape from the resonance, it is transferred to the SD zone ($q>30$ AU) or to the Centaur zone ($q<30$ AU) but it eventually switches to those population, due to the dynamical influence of Neptune.  The densest zone in the  orbital element space of escaped plutinos corresponds to the ranges $30 < a < 100$ AU and $5^{\\circ} < i < 40^{\\circ}$ and perihelions near the orbit of Neptune. When escaped plutinos are transferred to the SD they are quickly locked into a mean motion resonance with Neptune (similar to the behavior of SDOs analyzed by Fern\\'andez et al. (\\cite{Fernandez04}) and Gallardo (\\cite{Gallardo06}). In the Centaur zone (this is the zone of $q < 30$ AU ) the distribution of escaped plutinos is similar to that of SDOs in the Centaur zone obtained by Di Sisto \\& Brunini (\\cite{Disisto07}).  The orbital evolution of escaped plutinos in the Centaur zone can be grouped into the four dynamical classes proposed by Di Sisto \\& Brunini  (\\cite{Disisto07}). There are more particles that have the dynamical behavior of the second class and it is notable the great frequency of the presence of Kozai resonances in all the four classes. There are also several mean motion resonances densely populated in the ranges of $30 < a < 50$ AU.  The escaped plutinos have a mean lifetime in the Centaur zone of $  108$ Myr, greater than that of Centaurs from SD of $  72$ Myr . Escaped-plutinos live more time than SDOs in the greater-perihelion Centaur zone, causing a slower diffusion to the inner Solar System of escaped-plutino orbits than of SDOs orbits.  The present rate of injection of plutinos with radius greater than 1 km to the Centaur zone is between 1.6 to 16 plutinos every 100 years and the number of plutino-Centaurs with radius greater than 1 km would be between $1.8 \\times 10^{6} - 1.8 \\times 10^{7}$. Both, the rate of injection and the number of Centaurs from plutinos are much less than the contribution from the SD obtained by Di Sisto \\& Brunini  (\\cite{Disisto07}). Then, plutinos would represent a secondary source of Centaurs and their contribution would be a fraction of less than  $6 \\%$ of the total Centaur population.  \\vspace*{1cm}  \\noindent{\\bf Acknowledgments:} We thank  Matthew S. Tiscareno who, as referee, made valuable comments that helped to improve this manuscript."
    },
    "1005/1005.1865_arXiv.txt": {
        "abstract": "{} {We study the structure of the medium surrounding sites of high-mass star formation to determine the interrelation between the H~{\\sc ii} regions and the environment from which they were formed. The density distribution of the surrounding{s} is key in determining how the radiation of the newly formed stars { interacts with the surrounds in a way that allows it to be used as a star formation tracer}. } {We present new {{\\it Herschel/SPIRE}} 250, 350 and 500~$\\mu$m data of LHA 120-N44 and LHA 120-N63 in the LMC. We construct average spectral energy distributions (SEDs) for annuli centered on the IR bright part of the star formation sites. The annuli cover $\\sim$10$-$$\\sim$100 pc. We use a phenomenological dust model to fit these SEDs to derive the dust column densities, characterise the incident radiation field and the abundance of polycyclic aromatic hydrocarbon molecules. We see a factor 5 decrease in the radiation field {energy} density {as a function of radial distance} around N63. N44 does not show a systematic trend. We construct a simple geometrical model to derive the 3-D density profile of the surroundings of these two regions.} {{{\\it Herschel/SPIRE}} data {have proven} very efficient in deriving the dust mass distribution. We find that the radiation field in the two sources behaves very differently. N63 is more or less spherically symmetric and the average radiation field drops with distance. N44 shows no systematic decrease of the radiation intensity which is probably due to the {inhomogeneity} of the surrounding molecular material and to the complex distribution of several star forming clusters in the region.} {} ",
        "introduction": "\\label{sec:introduction} High-mass star formation (SF) sites (hereafter HMSFSs) are the beacons by which we probe a large part of the physics of external galaxies. They generally represent the most important tracers of the properties of their host galaxies in terms of star formation rate (SFR) and general activity. The most frequent tracers of the star-formation activity generally use the fact that the abundant UV light coming from the hot, young stars is absorbed in the vicinity and reradiated in the form of line or continuum emission. This is true, for example, for H$\\alpha$ \\citep[e.g.][]{1998ARA&A..36..189K}, the aromatic emission bands in the mid-IR \\citep[e.g.][]{2007ApJ...666..870C} or the IR continuum due to solid-state materials \\citep[dust, e.g.][]{1986ApJ...303L..41S}. These tracers work relatively well and are used to characterise nearby star forming regions and star forming galaxies out to large redshifts, although interesting discrepancies have been noted for dwarf galaxies at low SFR \\citep[see][]{2009ApJ...706..599L}. One of the main assumptions that enters into the quantitative interpretation of these data is the geometry of the material surrounding the newly formed stars, in particular, where the UV light is being reprocessed. For example, if the UV photons escape from the ionised medium this may boost the aromatic feature strengths and strongly influence the lines originating from the surrounding photo-dissociation regions (PDRs). There are indications that this geometry in external galaxies may qualitatively and quantitatively differ from that observed in the Milky Way (MW) \\citep[e.g.][]{2006A&A...446..877M,2009A&A...508..645G}. One simple effect may be that at different metallicities the surrounding medium is more or less opaque and therefore the UV photons have a different mean free path. More complex scenarios are also discussed in the literature. For example, clumpiness of the molecular cloud may lead to small molecular cores surrounded by large PDRs. {\\it Herschel} with its unprecedented wavelength coverage and angular resolution at submillimeter (submm) wavelengths provides a unique opportunity to probe the cold interstellar medium (ISM) and sample the effects of the environment on the resulting SF tracers. In particular, it is well suited to trace the distribution of matter around HMSFSs and to map the way the UV radiation permeates and heats the surroundings. Here we present a study of two HMSFSs in the Large Magellanic Cloud (LMC) based on data taken in the HERITAGE program \\citep[PI. Meixner, see][]{meixner_special_issue}. ",
        "conclusions": "\\label{sec:disc--concl} Figs.~\\ref{figResults} and \\ref{geometry} show the power of {\\it Herschel} to determine the matter distribution around HMSFSs. We investigate the effect the new SPIRE constraints on the derived parameters and their uncertainties by also fitting the SEDs without using these SPIRE data. We find that the column density of dust is often very discrepant (by more than an order of magnitude) from the values derived using the SPIRE data. The derived $\\Sigma_{dust}$ profile for N44 (like Fig.~\\ref{figResults} top panel) is virtually constant at a value of 0.1~M$_{\\sun}$\\,arcsec$^{-2}$ due to the fact that the cooler dust, is not well traced by the MIPS data. For N63 the model over-predicts the derived dust masses for the annuli between 20 to 60 pc by a large factor ($\\sim$5). In these cases the model tends to find very small values of $U_\\mathrm{min}$. This artifact could perhaps be circumvented, when modelling those kind of regions, in cases where submm constraints are missing, by limiting the allowed range of $U_\\mathrm{min}$ in Eq.~(\\ref{eqn:dale}). The difficulty of constraining the radiation field parameters without the SPIRE bands is also reflected in the derived uncertainties. In particular the value of $U_\\mathrm{min}$ is ill-constrained which results in large, asymmetric error-bars on the derived column densities.  N44 and N63 exhibit strikingly different behaviour in the radiation intensity profile (Fig.~\\ref{figResults}{\\bf d}.) The lack of a systematic decrease of $<$U$>$ around N44 indicates that we are observing dust with a wide range of temperatures along each line-of-sight. The inner annuli in N44 are affected by the superbubble to the NE of the OB association, where high values of $<$U$>$ are expected. The low values of $<$U$>$ for such a luminous SF region may reflect clumpiness. The profile is clearly incompatible with a centrally illuminated optically thin irradiation profile. It is clear from Fig.~\\ref{figMaps} that the studied regions are not very spherically symmetric (azimuthally smooth). In particular, N44 harbours several clusters and the peak of the X-ray emission is located in a cavity, $\\sim$ 20$-$30~pc away from the TIR peak (see Fig.~\\ref{figMaps}). Measuring the azimuthally averaged properties smears out some of the characteristics. This smearing could have been the cause for the lack of trend seen in the average $U$ as seen by the dust (Fig.~\\ref{figResults}). We have verified that this small range of $<$U$>$ is not simply an artifact of this averaging or a wrong choice of center of the annuli by studying the parameters we derive pixel by pixel in the maps which makes no assumptions about the geometry. Indeed the highest $<$U$>$ in N44 is found close to the center we chose. Except for the very center all other values, with their scatter, are within the range as depicted in Fig.~\\ref{figResults}. We conclude that the choice of center does not dominate the lack of trend of N44 in the average radiation field. Thus this lack of trend reflects the true broad range of irradiation conditions along all lines of sight in N44 , which is an indication of the inhomogeneity of the ISM around N44. A simple dust model shows a deficit in PAHs toward the centers of these two regions. We find no evidence for a submm excess. We have used the observed dust column densities surrounding N44 and N63 to derive a 3-D model for these regions for the first time."
    },
    "1005/1005.4995_arXiv.txt": {
        "abstract": "Pulsars are believed to be magnetized neutron stars. Their surface magnetic field ranges from $10^8$ to $10^{12}$ G. On the other hand, the magnetars have surface magnetic field $10^{14}-10^{15}$ G. It is believed that at the center, the magnetic field may be higher than that at the surface. We study the effect of the magnetic field on the neutron star matter and hence on the mass-radius relation of neutron stars. We model the nuclear matter with the relativistic mean field approach considering the possibility of appearance of hyperons at higher density. We find that the effect of magnetic field on the matter of neutron stars and hence on the mass-radius relation is important when central magnetic field $\\ge 10^{17}$ G. Moreover, if the central field is of the order of $10^{19}$ G, then the matter becomes unstable which limits the maximum magnetic field at the center of magnetars. ",
        "introduction": "Soon after the discovery of radio pulsar \\cite{hewish}, the theoretical proposition of neutron star (NS) \\cite{bz} drew much attention. At the core of a NS, the matter, composed of neutrons, protons, electrons, and sometimes muons, is highly degenerate with density in the range of $10^{14}-10^{15}$ gm/cc. Naturally, the constituent particles therein interact via nuclear forces, forming a non-ideal nuclear fluid. However, the nature of nuclear forces at this high density and the internal composition, especially the composition of matter in the inner most core of the NS, are not well understood yet. There are many theoretical models of nuclear matter at high density describing different equations of state for the matter. Based on different equations of state, different mass-radius relations of NS are obtained. Consequently, only way to constrain the equations of state is comparing the theoretical results with observed properties of NS.  With the detection of pulsars, it was also observed that the pulse period slowly increases in a very regular manner. Consequently, pulsars were successfully modeled as magnetized rotating NSs \\cite{pacini,gold,og}. This model successfully describes the spin-down power of the Crab pulsar and other pulsars. This description requires a surface magnetic field $10^{11}-10^{13}$ G for the pulsars detected at the beginning. Later on, more observations broadened this range of the magnetic field with the discovery of millisecond pulsars \\cite{backer}. This class of pulsars has a much lower surface magnetic field of $10^8-10^{10}$ G \\cite{alpar}. In the recent past pulsars with a very high rate of change of period (period derivatives) have been observed \\cite{morriset,mcet}, implying a high surface field up to $10^{14}$ G. Another class of pulsars has been observed which is the accreting X-ray pulsar \\cite{parmar}. For this class, the surface magnetic field varies from $10^8$ to $10^{13}$ G.  Anomalous X-ray pulsars (AXPs) were discovered in early 80s, which show periods in a relatively narrow range of $5-11$ s and have large period derivatives \\cite{vg}. In early days, they were believed to be unusual accreting NSs, whereas, later surveys detected no companions for these candidates. The large value of period derivatives implies a surface magnetic field as high as $10^{14}-10^{15}$ G. This helps in explaining the X-ray luminosities, powered by magnetic field decay, without any accretion. However, until the discovery of soft $\\gamma$-ray repeaters (SGRs), such high values of surface magnetic field were not acceptable. SGRs show brief intense bursts of soft $\\gamma$-rays. From the peak luminosities of the bursts, SGRs are identified as NSs with very strong surface magnetic field ($\\sim 10^{14}-10^{15}$ G) \\cite{nature2,woodset}. Now both AXP and SGR are believed to be strongly magnetized NSs with different variations, which are popularly known as magnetars \\cite{vg,dt,usov,td9}.  The existence of a magnetar motivates to study the effects of strong magnetic field on NS properties. A strong magnetic field affects, the structure of a NS through its influence on the underlying metric \\cite{bbgn,cpl} and equation of state (EoS) through the Landau quantization of charged particles and then the interaction of magnetic moments of charged particles (and even the anomalous magnetic moment (AMM) of neutral particles) with the magnetic field. In the present paper, we plan to study the effect of magnetic field on the EoS of highly dense matter. Generally, the effect of magnetic field on the EoS of matter is significant when the magnetic field ${\\cal B}>10^{18}$ G. For the nuclear matter with a $n$-$p$-$e$ system, the effect of magnetic field was studied by several authors \\cite{cbs,yz,bpl1,czl,wmkksg}. However, the composition of the core of a NS is very uncertain as the density reaches $10-15$ times the normal matter density. At this high density, the quarks may be deconfined, making quark matter, or the hyperons may appear, making hyperonic matter. Another hypothesis is that the whole star is made of deconfined quark matter, which is known as strange star (SS). The effect of magnetic field on quark matter using the MIT bag model has been studied in different literature \\cite{chakra,gc,fmro}. There are other models of quark matter with phenomenological density dependent quark masses \\cite{fowler,chak,d98,lxl}. Broderick et al. \\cite{bpl2} studied the effect of strong magnetic field on hyperonic matter. In their study, the field strength does not depend on density. However, in realistic situation, it is believed that the field strength is higher at core than that at surface of a NS. Therefore, with radius, and hence with density, the field strength should vary. In the present work, we plan to study the influence of the strong magnetic field on the hyperonic matter and hence on the mass-radius relation of NSs with hyperonic matter in the core with a density dependent magnetic field.  The paper is organized as follows. In the next section we describe the model of hadronic matter at high densities in the presence of magnetic field. Subsequently, we discuss the numerical results in section \\ref{discuss}. Finally, we summarize the results in the last section. ",
        "conclusions": "\\label{discuss}  Nucleon-meson coupling constants are chosen in such a way that the nuclear matter properties can be reproduced as the binding energy $E/B=-16.3$ MeV, the saturation density $n_0=0.153$ fm$^{-3}$, the asymmetry energy coefficient $a_{asy}=32.5$ MeV and the incompressibility $K=240$ MeV, and are taken from previous work \\cite{glend8}. The hyperon-$\\omega$ coupling constants are determined from SU(6) symmetry of the quark model \\cite{dg,sdgmgs,sm}. The hyperon-$\\sigma$ coupling constants, however, are determined from the potential depth of hyperons in normal nuclear matter given by \\begin{equation} U_Y=-g_{\\sigma Y} \\sigma+g_{\\omega Y} \\omega_0, \\end{equation} where $Y$ corresponds to hyperon. We take the potential depth for $\\Lambda$-hyperon $U_\\Lambda=-30~MeV$, as obtained from the analysis of $\\Lambda$-hypernuclei \\cite{dg,fet}. From various experimental data \\cite{fet,ket} for $\\Xi$-hypernuclei, we obtain the potential depth for $\\Xi$-hyperon $U_\\Xi=-18~MeV$. Recent $\\Sigma$-hypernuclei data indicate that the corresponding potential depth to be repulsive \\cite{bet} and we take $U_\\Sigma= 30~MeV$. Note importantly that earlier work \\cite{bpl2} assumed hyperon-meson coupling constants to be fixed fractions of respective nucleon-meson coupling constants and for all the hyperons these fractions are same for a given meson. Hence, unlike our model, in that work the first appeared hyperon is $\\Sigma^+$, which does not appear at all (see Table \\ref{threshold}) in our model.  In a realistic situation, the magnetic field at the core may be higher than that at the surface and gradually may decrease towards the surface. For the magnetars, the maximum surface magnetic field is $10^{15}$ G. For this case, the core magnetic field may be of the order of $10^{17}-10^{18}$ G. Without any proper knowledge of magnetic field configuration inside a NS, following previous work \\cite{bcs} here we adopt a magnetic field profile \\begin{equation} {\\cal B}\\left(n_b/n_0\\right)={\\cal B}_s+{\\cal B}_c\\left\\{1-e^{-\\beta \\left( \\frac {n_b}{n_0} \\right)^\\gamma}\\right\\}, \\label{magprfl} \\end{equation} where $\\beta$ and $\\gamma$ are two parameters determining the magnetic field profile with given ${\\cal B}_s$ and ${\\cal B}_c$, and $n_b$ is the total baryon number density. We consider different values of field at the center, whereas surface field strength is taken to be ${\\cal B}_s=10^{15}$ G. We notice that the different values of ${\\cal B}_s$ do not affect the EoS considerably, and thus we do not show the results with other ${\\cal B}_s$. In the above parametrization, the magnetic field strength depends on the baryon number density. However, at each density the field is uniform and constant.  \\subsection{Equation of states} We compute the EoS self-consistently with the condition of beta-equilibrium and charge neutrality. It is found that, as described by the Figures given below, the effect of magnetic field is important for ${\\cal B}_c \\ge 10^{17}$ G. However, for ${\\cal B}_c < 10^{18}$ G, the effect is significant only when the field reaches ${\\cal B}_c$ at a very low density, away from the center, and remains almost constant up to center. Therefore, we restrict our study with ${\\cal B}_c \\sim\\,10^{18}$ G. At lower densities, the matter is composed of only neutrons, protons and electrons. Hence, at the low density regime, the particles which are affected by the magnetic field are electrons and protons. Since the electrons are highly relativistic, electron Fermi momentum is very large compared to electron mass. Therefore, the number of occupied Landau levels by electrons is very large, eventhough the field strength under consideration is larger than the critical field strength of electron by several orders. On the other hand, the field strength under consideration is very less than the critical field strength of protons. Consequently, the number of occupied Landau levels by protons is also large.  As density increases, the heavier particles appear gradually. In addition, the magnetic field increases with the increase of density. As a result, the number of occupied Landau levels gradually decreases for every species. The threshold densities for muons and other hyperons to appear for ${\\cal B} _c=10^{18}$ G are given in Table \\ref{threshold}. The threshold densities for various species to appear do not differ from their respective values when the magnetic field is absent. \\begin{table}[t] \\begin{center} \\begin{tabular}[t]{|c|c|} \\hline \\multicolumn{2}{|c|}{Threshold densities} \\\\ \\hline $\\mu^-$ & 0.9 \\\\ \\hline $\\Lambda$ & 2.6 \\\\ \\hline $\\Xi^-$ & 3.1 \\\\ \\hline $\\Xi^0$ & 5.7 \\\\ \\hline \\end{tabular} \\caption{Threshold densities for muons and hyperons to appear in units of $n_0$.} \\label{threshold} \\end{center} \\end{table}  In Fig. \\ref{eos} we show the EoSs with and without magnetic field when the magnetic field profile corresponds to $\\beta =0.1$ and $\\gamma=1$ [see Eq. (\\ref{magprfl})]. We observe that the EoS becomes softer with the increase of ${\\cal B}_c$. Here we should mention that, when the field strength is high enough, the field energy and the field pressure are not negligible. In calculating the EoS, we thus add this contribution too which is necessary to construct the structure of NSs.   \\begin{figure}[t] \\begin{center} \\includegraphics[angle=-90,width=4.5 in]{p-ndiffb.eps} \\caption{Variation of $P$ as a function of normalized baryon number density. The solid curve is for without magnetic field. The dashed, dotted and dot-dashed curves correspond to ${\\cal B}_c=10^{18}$, $2\\times10^{18}$ and $3\\times 10^{18}$ G respectively.} \\label{eos} \\end{center} \\end{figure}  \\subsection{Mass-radius relations}  With the above EoSs, the structure of NSs with and without rotation is constructed. In this preliminary study of NS structure, we do not include anisotropic solutions based on the general relativity which is important in strong magnetic field. Without any rotation, the Einstein's equation leads to the Tolman-Oppenheimer-Volkoff (TOV) equation  \\begin{equation} \\frac{dP}{dr} = -\\frac{[\\varepsilon(r)+P(r)][M(r)+4\\pi r^3P(r)]} {r[r-2M(r)]}, \\end{equation} where $r$ is the distance from the center of the star and $M(r)$ the mass enclosed within $r$. The radius of a star ($R$) is the distance from the center at which the pressure becomes zero. For different central densities, stars with different masses and radii are formed. In Fig. \\ref{m-r} the mass-radius relations for static stars which correspond to EoSs described in Fig. \\ref{eos} are shown. Since the effect of magnetic field softens the EoS, in the presence of magnetic field the maximum mass is less than that for without magnetic field. Stars with rotation, as expected, have larger radii than their static counterpart. In Figs. \\ref{m-rr1} and \\ref{m-rr2} we show the corresponding mass-radius relations for rotating NSs. In both the cases, the maximum masses in presence of the magnetic field are less than that in absence of the magnetic field.  \\begin{figure}[t] \\centering \\includegraphics[angle=-90,width=4.5 in]{m-rst.eps} \\caption{Mass-radius relations without rotation. The solid curve is for without magnetic field. The dashed, dotted and dot-dashed curves correspond to ${\\cal B}_c=10^{18}$, $2\\times 10^{18}$ and $3\\times10^{18}$ G respectively.} \\label{m-r} \\end{figure}  \\begin{figure}[t] \\centering \\includegraphics[angle=-90,width=4.5 in]{m-rr2.eps} \\caption{Same as Fig. 2, except for rotating NSs with rotational frequency $\\Omega=2000$ s$^{-1}$.} \\label{m-rr1} \\end{figure}  \\begin{figure}[t] \\centering \\includegraphics[angle=-90,width=4.5 in]{m-rr4.eps} \\caption{Same as Fig. 2, except for rotating NSs with rotational frequency $\\Omega=4000$ s$^{-1}$.} \\label{m-rr2} \\end{figure}  \\subsection{Matter instability due to magnetic field} It has been noted that there is an onset of matter instability for certain sets of $\\beta$ and $\\gamma$ (which correspond to different magnetic field profiles) for every ${\\cal B}_c$. Now we explore EoSs for different magnetic field profiles to understand the instability. Figures \\ref{18eos}a and \\ref{318eos}a show the variations of $P$ at ${\\cal B}_c=10^{18}$ and $3 \\times 10^{18}$ G respectively, for different sets of $\\beta$ and $\\gamma$. For each $\\beta$, we take minimum $\\gamma$ to be $1$ for which ${\\cal B}$ varies very slowly with $n_b$ (see Figs. \\ref{18eos}b and \\ref{318eos}b). The maximum value of $\\gamma$ for every $\\beta$ has been chosen in such a way that beyond a certain value of $n_b$, $P$ ceases to increase (and decreases) with the further increase of $n_b$. This implies that the matter becomes unstable after that value of $n_b$. Keeping $\\beta$ constant if $\\gamma$ decreases, then instability disappears at some value of $\\gamma$ and the EoS tends to become that for without magnetic field. The point to note is that if $\\beta$ decreases, the instability occurs at larger values of $\\gamma$ and $n_b$, which is evident from Figs. \\ref{18eos}a and \\ref{318eos}a.  The instability occurs because the magnetic pressure contributes negatively to the matter pressure as reflected from the Eqs. (\\ref{enden}) and (\\ref{press}). Since the field strength increases with the increase of density, more negative contribution is added to the pressure, and finally, at a certain density, the pressure ceases to increase with the increase of the density. This effect has not been discussed in any earlier literature to the best of our knowledge. This is consistent with the result of softening the EoS with the increase of magnetic field. Since the softening of the EoS implies that with the energy density the rate of increase of pressure decreases, above a certain value of magnetic field, the rate becomes zero, and subsequently becomes negative. That is why, above a certain value of magnetic field the matter becomes unstable. This constrains the value of central magnetic field of a NS having high surface magnetic field, e.g. magnetar. In this context, it is worth mentioning that the inclusion of AMM effect \\cite {bpl2} stiffens the EoS, which dominates over the effect due to Landau quantization on EoS when field strength is $>5\\times10^{18}$ G. This stiffening is only possible when the field pressure is not considered in obtaining the EoS. However, the inclusion of field pressure further softens the EoS, as discussed above.  When the field strength is $\\sim 5\\times 10^{19}$ G, the deviation in EoSs between the cases with and without AMM effect is significant throughout the density range \\cite{bpl2}. However, this deviation is much small compared to that due to the extra effect arised from field pressure. Hence, at such a strong magnetic field, AMM effect is unable to stiffen the EoS. It is, however, seen that when the field strength is near $5\\times10^{18}$ G, AMM effect is insignificant at high density, but significant at low density. However, one should also remember that this result is obtained by considering constant field strength over the whole range of density. Thus, in the realistic case, when the field strength is several orders lower than $10^{18}$ G at low density, i.e., near the surface, AMM effect does not play any role in modifying the EoS. Therefore, we do not consider the AMM effect in our calculation while ${\\cal B}_c\\leq10^{19}$ G.  \\begin{figure}[t] \\begin{center} \\includegraphics[angle=-90,width=4.5 in]{18paper.eps} \\caption{Variation of (a) $P$, and (b) $B$ with normalized baryon number density for different magnetic field profiles with ${\\cal B}_c=10^{18}$ G. The solid curve is for without magnetic field. The sets of dashed, dotted, dot-dashed and short-dashed curves are for $\\beta=0.1$, $0.01$, $0.001$ and $0.0001$ respectively. In each set (for each $\\beta$), the upper curve in (a) is for $\\gamma=1$. The values of $\\gamma$ for the lower curve in sets in (a) are given in Table \\ref{gamma}. In (b), for each $\\beta$, the curves which reach ${\\cal B}_c$ at a comparatively lower $n_b$ correspond to maximum $\\gamma$ cases.} \\label{18eos} \\end{center} \\end{figure}  \\begin{figure}[t] \\begin{center} \\includegraphics[angle=-90,width=4.5 in]{318paper.eps} \\caption{Same as Fig. \\ref{18eos} except for ${\\cal B}_c=3\\times10^{18}$ G.} \\label{318eos} \\end{center} \\end{figure}   \\begin{table}[t] \\begin{center} \\begin{tabular}[h]{|c|c|c|} \\hline $\\beta$ & \\multicolumn{2}{c|}{$\\gamma$} \\\\ \\cline {2-3} & ${\\cal B}_c=10^{18}$ G & ${\\cal B}_c=3 \\times 10^{18}$ G \\\\ \\hline 0.1 & 4.0 & 1.4 \\\\ 0.01 & 6.0 & 2.6 \\\\ 0.001 & 6.0 & 3.5 \\\\ 0.0001 & 7.5 & 4.4 \\\\ \\hline \\end{tabular} \\caption{$\\gamma$-s for different lower curves shown in Figs. \\ref{18eos}a and \\ref{318eos}a for different $\\beta$-s.} \\label{gamma} \\end{center} \\end{table}  In Fig. \\ref{m-rallst}, the mass-radius relations for different ${\\cal B}_c$ as well as that without magnetic field have been shown. We notice that for a fixed magnetic field profile, e.g. $\\beta=0.1$ and $\\gamma=1$, the maximum mass decreases as ${\\cal B}_c$ increases from $10^{18}$ G to $4\\times10^{18}$ G. Now as ${\\cal B}_c$ increases further from $4\\times10^{18}$ G, with the same magnetic field profile (same $\\beta$ and $\\gamma$), $P$ starts decreasing with the increase of $n_b$ at a very low $n_b$, leading to unstable matter. For ${\\cal B}_c=5\\times10^{18}$ G, we obtain stable matter EoS for $\\beta=0.01$ and $\\gamma=2$, which deviates maximum from that without magnetic field, as shown in Fig. \\ref{518eos}. It is found that the maximum deviation of mass-radius relation for ${\\cal B}_c=5 \\times 10^{18}$ G with respect to that without magnetic field is less than that deviation for ${\\cal B}_c=4 \\times 10^{18}$ G (see Fig. \\ref{m-rallst}). Now if ${\\cal B}_c$ increases further to $10^{19}$ G, we note that (see Fig. \\ref{19eos}) with the above said profile ($\\beta=0.01$ and $\\gamma=2$) the matter becomes unstable at a very low density. We have stable matter up to a reasonably high density with ${\\cal B}_c=10^{19}$ G for $\\beta=0.001$ and $\\gamma=3$ with maximum deviation in EoS from that without magnetic field. However, Fig. \\ref{m-rallst} shows that in this case the deviation of EoS from that without magnetic field case is even lesser. Hence we can conclude that the effect of magnetic field is significant when ${\\cal B}_c$ is around $4\\times10^{18}$ G depending on the magnetic field profile. Beyond that, the magnetic field profiles needed to create significant effect of magnetic field on the mass-radius relation make the matter unstable.  \\begin{figure}[t] \\begin{center} \\includegraphics[angle=-90,width=4.5 in]{m-rallst.eps} \\caption{Mass-radius relations for static star with different values of ${\\cal B}_c$. The solid curve is for without magnetic field. The dashed curves are for ${\\cal B}_c=10^{18}$, $3\\times 10^{18}$ and $4\\times10^{18}$ G respectively from top to bottom for $\\beta =0.1$ and $\\gamma=1$. The dotted curve corresponds to ${\\cal B}_c=5\\times10^{18}$ G for $\\beta=0.01$ and $\\gamma=2$, and the dot-dashed curve is for ${\\cal B}_c=10^{19}$ G when $\\beta=0.001$ and $\\gamma=3$.} \\label{m-rallst} \\end{center} \\end{figure}  \\begin{figure}[t] \\begin{center} \\includegraphics[angle=-90,width=4.5 in]{518eos.eps} \\caption{$P$ as a function of normalized baryon number density with ${\\cal B}_c=5\\times10^{18}$ G for different magnetic field profiles. The solid curve is for without magnetic field, the dashed curve for $\\beta=0.1$ and $\\gamma=1$, the dotted curve for $\\beta=0.01$ and $\\gamma=3$, the dot-dashed curve for $\\beta=0.01$ and $\\gamma=2$ and the short-dashed curve for $\\beta=0.001$ and $\\gamma=4$.} \\label{518eos} \\end{center} \\end{figure}  \\begin{figure}[t] \\begin{center} \\includegraphics[angle=-90,width=4.5 in]{19eos.eps} \\caption{$P$ as a function of normalized baryon number density with ${\\cal B}_c=10^{19}$ G for different magnetic field profiles. The solid curve is for without magnetic field. The dashed curve is for $\\beta=0.01$ and $\\gamma=2$, the dotted curve for $\\beta=0.01$ and $\\gamma=1$ and the dot-dashed curve for $\\beta =0.001$ and $\\gamma=3$.} \\label{19eos} \\end{center} \\end{figure}  As we have mentioned earlier, the effect of magnetic field on quark matter using the MIT bag model has been studied in many earlier work. Therefore, further it will be interesting to see the effect of magnetic field on other models, e.g. of Dey et al. \\cite{d98}, which is in progress."
    },
    "1005/1005.2526_arXiv.txt": {
        "abstract": "{ The goal of local helioseismology is to elicit three-dimensional information about the sub-surface (or far-side) structure and dynamics of the Sun from observations of the helioseismic wave field at the surface. The physical quantities of interest include flows, sound-speed deviations and magnetic fields.  However, strong surface magnetic fields induce large perturbations to the waves making inversions difficult to interpret. The purpose of this paper is to outline the methods of analysis used in local helioseismology, review discoveries associated with the magnetic Sun made using local helioseismology from the past three years,  and highlight the efforts towards imaging the interior in the presence of strong magnetic fields. } ",
        "introduction": "Helio- and astero-seismology provide the means to ``see\" beneath the surface of the Sun or star, and to gain an understanding of the physics of stellar interiors. The analysis of the frequencies of the global modes of oscillation allows us to study the internal stellar structure only as a function of radius and unsigned latitude. With high spatial resolution observations of the velocity signal on the Sun's surface, local helioseismology is able to image the three dimensional subsurface structure of the Sun and image magnetic activity on the far-side. Local helioseismology aims to interpret the full wave field observed at the solar surface, not only the frequencies of the normal modes. We restrict our attention, here, to local helioseismology.  The Doppler signal, which is the line-of-sight velocity of the solar surface, is the fundamental observable. Full-disk high-quality data is obtained from instruments including the Michelson Doppler Image (MDI) onboard the Solar and Heliospheric Observatory (SOHO) \\citep{MDI95}, the Global Oscillation Network Group  (GONG) \\citep{Harveyetal1996}, and the Helioseismic and Magnetic Imager (HMI) onboard the recently launched Solar Dynamics Observatory (SDO). MDI also runs higher-resolution campaigns on smaller fields-of-view. The Magneto-Optical filter at Two Heights \\citep{Finsterle2004} observes simultaneously at two heights in the atmosphere allowing an analysis of the vertical propagation of waves.  Various local helioseismic techniques exist to analyse the oscillations. The main ones are time-distance analysis \\citep{DJHP93}, ring diagram analysis \\citep{Hill1988} and acoustic imaging/holography \\citep{CCL97,LB97}. Ring-diagram analysis is most closely related to global mode helioseismology. By analysing frequency shifts over small regions of the solar surface, the direction and amplitude of subsurface flows can be determined. Time-distance analysis computes the travel time of a wave packet travelling between two points on the solar surface. Helioseismic holography (acoustic imaging) reconstructs the subsurface wave fields by propagating waves either forward or backward in time. Other methods include Fourier-Hankel analysis \\citep{BDL87}, designed specifically to analyse waves surrounding a sunspot. \\cite{Woodard2009} used direct modelling to interpret the wave correlations in wavevector-frequency space.  The forward problem is then to compute the expected seismic observations from a particular model of the Sun. The inverse problem is to infer the internal properties of the Sun from the seismic observables. Local helioseismology is the only way to elicit the three-dimensional structure of the solar interior, including structures and flows on a range of scales.   In this paper we focus on the local helioseismology of sunspots,  active regions and large scale flows varying with the solar cycle. There are particular problems associated with doing helioseismology in the immediate vicinity of strong surface magnetic fields. The forward problem used to infer the subsurface properties has previously been based exclusively on solving hydrodynamic equations in the absence of a magnetic field. Additionally, the actual inversion method assumes that any perturbation to the waves is small, however it has been demonstrated that strong surface magnetic fields cause large perturbations to the waves. Knowing the subsurface structure of sunspots and active regions will constrain sunspot models and hopefully determine their deeper structure. The immediate need now is to quantitatively model wave propagation through magnetic fields and, subsequently, account for the effects when inverting for the subsurface structure.  An important question in solar physics is to understand the mechanics of the solar dynamo.  Local helioseismology has identified large scale flow variations that appear to be intimately connected to the surface activity. Understanding such connections will help to decipher the larger scale dynamics of the magnetic Sun. Other goals of local helioseismology are to measure the meridional flow and changes at the base of the convection zone, which would help constrain models of the dynamo.  This short review will cover the basic techniques used in local helioseismology  in Section~\\ref{at}. How the interior structure is subsequently inferred is described in Section~\\ref{inv}. Section~\\ref{obs} will then cover some of the results, over the past three years, from local helioseismic analysis tied to solar activity. A concentrated effort towards modelling waves propagating through magnetic fields to  understand the observations will be reviewed in Section~\\ref{mod}. Recent reviews of local helioseismology include \\cite{KHH06,GT07,Birch2008}. For more extensive reviews of local helioseismology see, for example, \\cite{GBS10}. ",
        "conclusions": "Further confirmation of the surface effects of strong magnetic fields have  been found in observations, casting further doubt on over-simplified inversions for wave-speed and flows below magnetic active regions and sunspots. The effects of filtering, source suppression  and mode conversion of the waves in strong surface magnetic fields have yet to be fully quantified for each of the analysis techniques. In turn, a concerted effort is underway to quantitatively simulate the seismic signature of the Sun in the presence of magnetic fields to explore these effects. The simulations are still being fully analysed and in the near future a quantitative measure of the effects may be used to infer the correct subsurface properties of a sunspot.  Future efforts in local helioseismology will focus on deeper inversions, particularly of the base of the convection zone, and steps towards this  have been taken using time-distance techniques \\citep{ZHKM09}.  A challenging goal is to detect the emergence of magnetic flux before it is observed at the surface. Some progress has been made towards this \\citep[e.g.][]{KMHH08,KHH09,HKZM10}, but strong definitive signatures have not been observed. Future studies of the effects of magnetic flux emergence will require systematic helioseismic analysis of many emerging active regions and realistic simulations.  The HMI instrument onboard the recently launched  \\linebreak SDO  very soon provide increased resolution observations of the seismic Sun and provide accompanying vector magnetograms on a regular basis. Thanks to the improved spatial resolution we anticipate being able to do local helioseismology closer to the limb (and poles)."
    },
    "1005/1005.5076_arXiv.txt": {
        "abstract": "{We point out some of the outstanding challenges for embedding inflationary cosmology within string theory studying the process of reheating for models where the inflaton is a closed string mode parameterising the size of an internal cycle of the compactification manifold. A realistic model of inflation must explain the tiny perturbations in the cosmic microwave background radiation and also how to excite the ordinary matter degrees of freedom after inflation, required for the success of Big Bang Nucleosynthesis. We study these issues focusing on two promising inflationary models embedded in LARGE volume type IIB flux compactifications. We show that phenomenological requirements and consistency of the effective field theory treatment imply the presence at low energies of a \\textit{hidden sector} together with a {\\it visible sector}, where the Minimal Supersymmetric Standard Model fields are residing. A detailed calculation of the inflaton coupling to the fields of the hidden sector, visible sector, and moduli sector, reveals that the inflaton fails to excite primarily the visible sector fields, instead hidden sector fields are excited copiously after the end of inflation. This sets severe constraints on hidden sector model building where the most promising scenario emerges as a pure $N=1$ SYM theory, forbidding the kinematical decay of the inflaton to the hidden sector. In this case it is possible to reheat the Universe with the visible degrees of freedom even though in some cases we discover a new tension between TeV scale SUSY and reheating on top of the well-known tension between TeV scale SUSY and inflation.}  \\preprint{DESY 10-072}  \\begin{document}  \\tableofcontents  \\bigskip ",
        "introduction": "Primordial inflation has been a very successful paradigm which can explain the observed temperature anisotropy in the cosmic microwave background (CMB) radiation~\\cite{WMAP7} (for a recent review see~\\cite{RM}). Any successful model of inflation must also explain how to excite the Standard Model (SM) quarks and leptons after the end of inflation through the process of reheating (for a review see~\\cite{Reheat-rev}) required for the success of Big Bang Nucleosynthesis~\\cite{BBN}. In this regard only visible sector models of inflation can be considered to be safe where the inflaton carries the SM charges~\\cite{MSSM-inf}. Such models are primarily based on the Minimal Supersymmetric Standard Model (MSSM), where supersymmetry (SUSY) helps in obtaining a {\\it cosmologically flat} potential via $D$-flat directions (for a review on MSSM flat directions see~\\cite{MSSM-rev}). The advantage of visible sector models  embedded within the MSSM is that it can also reproduce the required dark matter abundance necessary for the structure formation~\\cite{ADM}.  On the contrary there exists a plethora of models of inflation where the inflaton belongs to the hidden sector since it is a SM gauge singlet. Typically these models can explain the temperature anisotropy in the CMB~\\cite{RM}, but fail to explain how the inflaton energy density gets transferred primarily to visible sector degrees of freedom (\\textit{dof}), and not to hidden sector ones. Note that \\textit{a priori} a gauge singlet inflaton has no preference to either the visible or the hidden sector. Therefore, it is a challenge to construct a model of inflation where the inflaton belongs to the hidden sector and motivate the inflaton couplings to hidden and visible \\textit{dof}.  It is typical to encounter such a scenario in string theory which provides many closed string modes that are good inflaton candidates. In particular, the low energy effective action of string compactifications on Calabi-Yau three-folds typically has a large number of uncharged massless scalar fields with a flat potential, called moduli. These moduli can be the ideal candidates to drive inflation. Since their vacuum expectation value (VEV) determines various parameters (like masses and coupling constants) of the four-dimensional effective theory, it is important to stabilise them by providing masses to these moduli. Otherwise, the presence of these massless scalars with effective gravitational coupling would mediate unobserved long range fifth forces. It is well known by now that nonzero background fluxes induce potentials for some of the moduli~\\cite{gvw,drs,gkp}. However these fluxes severely backreact on the Calabi-Yau geometry generically modifying the internal space to a manifold with a more complicated structure. This is not the case for type IIB compactifications where the fluxes induce just a warp factor. For this reason, moduli stabilisation is best understood in the context of type IIB which is therefore well suited for the discussion of phenomenology and cosmology. However the stabilisation of all the geometric moduli requires to take into account various perturbative and non-perturbative effects \\cite{kklt, LVS, GenAnalofLVS}.  One of the nice properties of type IIB flux compactifications is that global (or bulk) issues decouple from local (or brane) issues. The former are model independent issues which depend on moduli stabilisation and some of them are: (a)~Different hierarchical scales,~(b)~SUSY breaking,~(c)~Soft SUSY breaking mass terms,~(d)~Cosmological constant, and ~(e)~Inflation. On the other hand, the local issues depend on the particular local $D$-brane construction and some of them are: (a)~ Gauge group,~(b)~Chiral spectrum,~(c)~Yukawa couplings,~(d)~Gauge coupling unification, ~(e)~Proton stability, and ~(f)~Baryogenesis.  An excellent example of type IIB compactifications with stabilised moduli is given by the LARGE Volume Scenario (LVS) originally proposed in \\cite{LVS}. This is very appealing for particle physics phenomenology and cosmology. In these compactifications $\\alpha'$ and $g_s$ corrections can be combined with non-perturbative effects in order to generate a potential for all the K\\\"{a}hler moduli \\cite{GenAnalofLVS}, whereas the background fluxes induce a potential for the dilaton and the complex structure moduli. Unlike the KKLT set-up \\cite{kklt}, the moduli stabilisation is performed without fine tuning of the values of the internal fluxes and the Calabi-Yau volume is fixed at an exponentially large value (in string units). As a consequence, one has a very reliable four-dimensional effective description, as well as a tool for the generation of phenomenologically desirable hierarchies.  There are two different choices to embed the visible sector within LVS: \\begin{enumerate} \\item {\\it Geometric regime}: the visible sector is built via magnetised intersecting $D7$ branes wrapping a blow-up 4-cycle which is stabilised at a value larger than the string scale~\\cite{GeomRegime};  \\item {\\it Quiver locus}: the visible sector is built via fractional $D$-branes at the singularity obtained by shrinking down a blow-up mode~\\cite{quiver}. \\end{enumerate} Both of them have different phenomenology, the scales and hierarchies are different and we will briefly review them in the following section.  One of the interesting features for all these setups is that at low energies they always come with a hidden and a visible sector. Therefore these models will serve as a perfect playground to study the inflaton couplings to both the \\textit{dof}. Furthermore, the two sectors do not have any direct coupling among themselves, except indirectly via the closed string sector. In this respect the two sectors are isolated from each other, nevertheless their total energy density will govern the dynamics of the Universe. It is then crucial to find an answer to each of the following questions: how does the inflaton couple to each sector? how does perturbative reheating take place? is it at all possible to excite predominantly the MSSM \\textit{dof}? Thankfully, these questions can be answered within this setup which involves an interesting interplay between local and global issues. We shall stress that reheating can set severe constraints both on the hidden sector physics and the scale of inflation which should be compatible with CMB predictions, successful reheating, and TeV scale SUSY\\footnote{Reheating and thermalisation have been studied previously in the context of closed string excitations in brane inflation~\\cite{Many-reheat}, in the case of warped geometry~\\cite{Warped-reheat}, and in closed string models of inflation~\\cite{Green-reheat, Preheating}. In all these cases, no clear distinction has been made between hidden and visible \\text{dof}. In this respect such studies are interesting but leave little impact on how the MSSM \\textit{dof} are excited, whose answer is relevant for the success of BBN.}.  In order to study these issues, we shall focus on two very promising models of string inflation embedded in LVS, where the inflaton is a closed string mode, more precisely a K\\\"{a}hler modulus parameterising the size of an internal 4-cycle: (1)~{\\it Blow-up inflation} (BI) \\cite{kahlerinfl}: in this case the inflaton is the size of a blow-up mode yielding a small field inflationary model which is in good agreement with current observational data; and (2)~{\\it Fibre inflation} (FI) \\cite{fiberinfl}: in this case the inflaton is the size of a K3 fibre over a $\\mathbb{C}P^1$ base producing a large field inflationary model which can be a potential candidate if primordial gravity waves are discovered. These models represent some of the few known examples where inflation can be achieved without fine-tuning any parameter in the potential. In fact it is the typical no-scale structure of the potential that allows to solve the $\\eta$-problem in a natural way.  In spite of the successes of these models we shall discover two outstanding challenges for closed string inflationary scenarios: \\begin{enumerate} \\item Most of the inflaton energy gets dumped to hidden sector \\textit{dof}. This will lead to overproduction of hidden sector dark matter if there exists a stable species. This problem seems to rule out all the inflationary models unless the hidden sector consists just of an $N=1$ pure SYM theory. This severe constraint on hidden sector model building can be relaxed only for the geometric regime case with the inflaton wrapped by the visible sector $D7$-stack.  \\item Incompatibility between TeV scale SUSY and reheating on top of the well-know tension between inflation and TeV scale SUSY \\cite{InfvsSUSY}: assuming that the density perturbations can be generated in a non-standard way by a second curvaton or modulating field \\cite{BCGQTZ}, one could set the scale of the scalar potential in order to get inflation, the correct amount of CMB temperature fluctuations and TeV scale SUSY at the same time. However it might still be impossible to achieve an efficient reheating of the visible sector due to the fact that the inflaton might decay after BBN. This problem seems to rule out all the FI models (with the possible exception of FI at the quiver locus) and BI in the geometric regime with the inflaton not wrapped by the visible sector. The only models which are left over are BI in the geometric regime with the inflaton wrapped by the visible sector or BI at the quiver locus. It is important to stress that both these models require fine-tuning to have a successful inflationary scenario. \\end{enumerate}  This paper is organised as follows. In section 2, we give a brief review of LVS stressing the different features of the geometric regime and the quiver locus case. We then discuss the dynamics of BI and FI presenting their predictions for the cosmological observables. We finally briefly describe the process of perturbative and non-perturbative reheating and thermalisation. In section 3, we discuss the dynamics of the hidden sector and the constraints on hidden sector \\textit{dof}. In section 4, we describe the inflaton couplings to both hidden and visible \\textit{dof} estimating the maximal reheating temperature in the approximation of sudden thermalisation of the visible sector for BI and FI, respectively. In section 5, we discuss our main results and conclude. In appendix \\ref{appendice}, we present a detailed derivation of the moduli canonical normalisation, mass spectrum, couplings and decay rates to any particle present in our models both for BI and FI. More precisely, we have derived the moduli couplings to fermions/sfermions, Higgs/Higgsinos, gauge bosons/gauginos of the visible and hidden sector, and the moduli self couplings. ",
        "conclusions": "Once the inflaton has completely decayed into visible sector \\textit{dof}, the system evolves to a complete thermalisation and the thermal bath generates a moduli-dependent finite temperature scalar potential\\footnote{Note that within the MSSM, as it is the case for many scenarios, the final reheating temperature can be as low as $1$ TeV.}. The behaviour of this potential for an arbitrary LVS has been studied in \\cite{LVSatFiniteT} where it was found that the finite temperature corrections do not develop any new minimum, and so they cannot induce any phase transition, but they can still lead the system to a dangerous decompactification limit. The computation of the decompactification temperature for LVS gives \\cite{LVSatFiniteT}: \\be T_{max}\\simeq\\frac{M_P}{\\vo^{3/4}}, \\label{Tmax} \\ee and this is valid for all the different scenarios that we have studied here. This temperature sets the maximal temperature of our Universe, and so one has to make sure that $T_{RH}^{max}<T_{max}$. It is reassuring to notice that in all our cases the reheating temperature is always suppressed with respect to the decompactification temperature (\\ref{Tmax}) by inverse powers of the overall volume. Only in the case of BI with the inflaton wrapped by the visible sector, the ratio $R\\equiv T_{RH}^{max}/T_{max}\\simeq \\vo^{-1/4}$ is not particularly small. Hence it is important to compute the exact coefficient in from of the $\\vo^{-1/4}$ factor. This has already been done in \\cite{LVSatFiniteT} and the final result is: \\be R\\equiv \\frac{T_{RH}^{max}}{T_{max}}= c\\frac{\\left(\\ln\\vo\\right)^{3/2}}{\\vo^{1/4}}\\text{ \\ with \\ } c\\simeq \\frac{W_0\\, e^{K_{cs}/2}}{2\\langle\\tau_2\\rangle^{11/8}}, \\ee For $\\vo\\simeq 10^{6- 7}$, the previous ratio just depends on the $c$ coefficient: $R\\simeq c$. However for natural values $W_0\\simeq K_{cs}\\simeq \\mathcal{O}(1)$ and $\\langle\\tau_2\\rangle\\simeq \\mathcal{O}(10)$, $c\\simeq 0.1- 0.01$, and so we correctly have $R < 1$. In the fine-tuned cases when $R$ becomes dangerously of the order one, larger values of the volume seem to be favoured. However for $\\vo>10^7$, the inflaton fluctuations are not able anymore to generate the right amount of density perturbations, which could be, on the other hand, produced by some non-standard mechanism, like a curvaton or a modulating field \\cite{BCGQTZ}.  \\medskip  Let us now summarise the main results we have found in this paper: \\begin{itemize} \\item In any LVS hidden sectors are always present along with the visible sector.  \\item The ubiquitous presence of a hidden sector wrapping the inflaton 4-cycle implies that BI needs fine-tuning whereas FI does not.  \\item The hidden sector and the visible sector are not directly coupled, and furthermore the inflaton coupling to visible sector \\textit{dof} can never be made stronger than its coupling to hidden sector \\textit{dof}.  \\item Two generic problems arise: \\begin{enumerate} \\item Inflaton energy dumping to hidden, instead of visible, sector \\textit{dof}. As such this is not a problem provided the hidden sector \\textit{dof} are relativistic and there exists a curvaton mechanism which is responsible for creating the perturbations and the visible sector matter. Such a scenario was considered in past~\\cite{EKM}, where the MSSM Higgs would dominate the energy density during its oscillations and would create all the visible \\textit{dof}. It would also be interesting to investigate if this problem could be solved by building a curvaton model along the lines of \\cite{BCGQTZ} where the curvaton dominates the energy density.  \\item Overproduction of hidden sector dark matter if there exists a massive long lived species, otherwise there would be too many new relativistic \\textit{dof}, which would spoil the BBN predictions. \\end{enumerate}  \\item A possible solution could be to render the inflaton decay to hidden sector \\textit{dof} kinematically forbidden. This sets severe constraints on hidden sector model building leaving as the most promising possibility just a pure $N=1$ SYM theory. This is always the case for each hidden sector at the quiver locus and in the geometric regime if the inflaton is not wrapped by the visible sector. On the contrary, in the geometric regime case with the inflaton wrapped by the visible sector, the requirement of having just a pure SYM theory for the hidden sector not wrapping the inflaton 4-cycle can be relaxed.  \\item This constraint seems to be in contrast with the generic presence of hidden photons whose study received a lot of attention recently \\cite{Mark}. More precisely, the appearance of hidden photons living on small or large cycles seems to be compatible with reheating only for models in the geometric regime with the inflaton 4-cycle wrapped by the visible sector. All the other cases seem to be incompatible with the existence of hidden photons. \\end{itemize}  Assuming that these problems can be solved by an appropriate brane construction of the hidden sector, let us summarise our results for the maximal reheating temperature of the visible sector:  \\begin{figure}[ht] \\begin{center} \\begin{tabular}{c||c|c} & Geometric Regime & Quiver Locus \\\\ \\hline\\hline \\\\ & & \\vspace{-0.9cm}\\\\ & Inflaton not wrapped by the visible sector:  \\\\ Blow-up & $T_{RH}^{max}\\sim M_P \\vo^{-2}\\sim 10^{6- 8}$ GeV for $\\vo\\sim 10^{6- 7}$ & $T_{RH}^{max}\\sim M_P \\vo^{-5/2}\\sim 10^{2- 4}$ GeV \\\\ Inflation & Inflaton wrapped by the visible sector: & for $\\vo\\sim 10^{6- 7}$ \\\\ & $T_{RH}^{max}\\sim M_P \\vo^{-1}\\sim 10^{12- 13}$ GeV for $\\vo \\sim 10^{6- 7}$ \\\\ \\hline \\\\ & & \\vspace{-0.9cm}\\\\ & Inflaton not wrapped by the visible sector:  \\\\ Fibre & $T_{RH}^{max}\\sim M_P \\vo^{-17/6}\\sim 10^{6- 9}$ GeV for $\\vo\\sim 10^{3- 4}$ & $T_{RH}^{max}\\sim M_P \\vo^{-10/3}\\sim 10^{4- 7}$ GeV \\\\ Inflation & Inflaton wrapped by the visible sector: & for $\\vo\\sim 10^{3- 4}$ \\\\ & $T_{RH}^{max}\\sim M_P \\vo^{-5/2}\\sim 10^{8- 10}$ GeV for $\\vo\\sim 10^{3- 4}$ \\end{tabular}\\\\ \\end{center} \\end{figure}  We stress that only the model of BI with the visible sector at the quiver locus gives rise to a viable reheating with TeV scale SUSY at the same time. In all other cases $M_{soft}\\gg 1$ TeV. Let us now assume that the inflaton is only responsible to drive inflation but the density fluctuations are generated in a non-standard way by a second curvaton or modulating field \\cite{BCGQTZ}. In this way we still obtain inflation but the volume could be set so to obtain TeV scale SUSY. The resulting maximal reheating temperature in the visible sector would turn out to be:  \\begin{enumerate} \\item BI in the geometric regime with the inflaton not wrapped by the visible sector: for $\\vo\\simeq 10^{14}$ (as needed to obtain TeV scale SUSY in the geometric regime), we would obtain $T_{RH}^{max}\\simeq 10$ eV. This is in clear disagreement with the lower bound coming from BBN: $T_{RH}^{max}>1$ MeV. Thus we infer that this model does not seem to be compatible with TeV scale SUSY.  \\item BI in the geometric regime with the inflaton wrapped by the visible sector: for $\\vo\\simeq 10^{14}$, we would obtain $T_{RH}^{max}\\simeq 5\\cdot 10^{5}$ GeV which is not in contrast with any observation. Thus we infer that this model could be made compatible with TeV scale SUSY.  \\item FI in the geometric regime with the inflaton not wrapped by the visible sector: for $\\vo\\simeq 10^{14}$, and $c\\simeq 5\\cdot 10^{-2}$ we would obtain $T_{RH}^{max}\\simeq 10^{-14}$ eV, which is unrealistically small. Thus we infer that this model does not seem to be compatible with TeV scale SUSY.  \\item FI in the geometric regime with the inflaton wrapped by the visible sector: for $\\vo\\simeq 10^{14}$, and $c\\simeq 5\\cdot 10^{-2}$ we would obtain $T_{RH}^{max}\\simeq 10^{-8}$ eV, which is unrealistically small. Thus we infer that this model does not seem to be compatible with TeV scale SUSY. On top of this, we would obtain an unrealistically small value for the visible sector gauge coupling: $\\alpha^{-1}_{vis}=\\langle\\tau_1\\rangle=c\\,\\vo^{2/3}\\simeq 10^8$.  \\item FI at the quiver locus : for $\\vo\\simeq 10^6$ (as needed to obtain TeV scale SUSY at the quiver locus), we would obtain $T_{RH}^{max}\\simeq 1$ MeV which is of the same order of magnitude as $T_{BBN}$. Thus we infer that this model could be made compatible with TeV scale SUSY only if $T_{RH}^{max}$ turns out to be actually a bit higher than $T_{BBN}$ and the matter-antimatter asymmetry is generated in a non-thermal way. \\end{enumerate}  Thus, even though FI does not require any fine-tuning of the coefficient of the string loop corrections, it seems very difficult to obtain a viable reheating for these models which is also compatible with TeV scale SUSY. On the contrary, both BI at the quiver locus and in the geometric regime with the inflaton wrapped by the visible sector, seems to be more promising to achieve inflation, reheating and TeV scale SUSY at the same time.  We finally stress that our estimate of the reheating temperature was based on the approximation of sudden thermalisation of the visible sector \\textit{dof}. In reality it is more likely that the system slowly evolves to thermal equilibrium with an actual reheating temperature $T_{RH}< T_{RH}^{max}$. However we shall leave the detailed study of the thermalisation process for future investigation."
    },
    "1005/1005.4593_arXiv.txt": {
        "abstract": "We show that we can obtain a good fit to the present day stellar mass functions (MFs) of a large sample of young and old Galactic clusters in the range $0.1 - 10$\\,\\Msolar with a tapered power law distribution function with an exponential truncation of the form $dN/dm \\propto m^\\alpha \\, [1 - e^{-(m/m_c)^\\beta}]$. The average value of the power-law index $\\alpha$ is $\\sim -2$, that of $\\beta$ is $\\sim 2.5$, whereas the characteristic mass $m_c$ is in the range $0.1 - 0.8$\\,M$_\\odot$ and does not seem to vary in any systematic way with the present cluster parameters such as metal abundance, total cluster mass or central concentration. However, $m_c$ shows a remarkable correlation with the dynamical age of the cluster, namely $m_c/M_\\odot \\simeq 0.15 + 0.5 \\times \\tau_{\\rm dyn}^{3/4}$, where $\\tau_{\\rm dyn}$ is the dynamical age taken as the ratio of cluster age and dissolution time. { The small scatter seen around this correlation is consistent with the uncertainties on the  estimated value of $\\tau_{\\rm dyn}$.} We attribute the observed trend to the onset of mass segregation via two-body relaxation in a tidal environment, causing the preferential loss of low-mass stars from the cluster and hence a drift of the characteristic mass $m_c$ towards higher values. { If dynamical evolution is indeed at the origin of the observed trend, it would seem plausible that high-concentration globular clusters, now with median $m_c \\simeq 0.33$\\,M$_\\odot$, were born with a stellar MF very similar to that measured today in the youngest Galactic clusters and with a value of $m_c \\simeq 0.15$\\,M$_\\odot$. This hypothesis is consistent with the absence of a turn-over in the MF of the Galactic bulge down to the observational limit at $\\sim 0.2$\\,M$_\\odot$ and, if correct, it would carry the implication that the characteristic mass is not set by the thermal Jeans mass of the cloud.} ",
        "introduction": "There is general consensus that stars do not form in isolation but rather from the fragmentation of molecular clouds that leads to star clusters (e.g. Lada \\& Lada 2003; Elmegreen 2010). This makes stellar clusters ideal places to study the properties of star formation and its outcome, the stellar initial mass function (IMF). However, it is also equally well established today that a large majority of stars, in our Galaxy and elsewhere, are not in clusters but in the field, since clusters disrupt over time (see e.g. Gieles 2010). Therefore, any attempt to set constraints on the star formation mechanisms from the analysis of the present day stellar mass function (MF) of but the youngest clusters cannot ignore the consequences of their dynamical evolution.  This issue becomes particularly important if we want to compare the results of star formation in physically different environments and with various ages. The obvious example is addressing differences in the way low-mass stars ($< 1$\\,M$_\\odot$) formed in globular clusters (GCs), at $z \\simeq 5$ or about 12\\,Gyr ago, and in young clusters (YCs) in the local universe. While in both cases the raw data show a broad plateau in the mass distribution, suggesting a characteristic mass of order a few tenths of M$_\\odot$ (e.g. Elmegreen et al. 2008), until the effects of dynamical evolution are properly understood, no meaningful conclusion can be drawn either on the uniformity of the  star formation process over time or on the possible role of the  environment. In order to cover a wide range of metallicities and initial densities one is forced to compare stellar systems that formed a Hubble time apart, such as the YCs and GCs mentioned above. Therefore, the universality of  the IMF over time and  location remains as yet an unresolved issue.  On the other hand, even though the two-body relaxation process that governs the cluster's dynamical evolution and that leads to preferential loss of low-mass stars is today rather well understood (e.g. Spitzer 1987), it is in general not possible to roll back the effects of dynamics and derive the IMF from the present day MF, since we  cannot trace back the trajectories of stars that have escaped the cluster. It is however possible and statistically meaningful to study the evolution of the stellar MF on a global scale by looking at the differences between clusters at different evolutionary stages. This requires a homogeneous sample of high quality observations of Galactic (open and globular) clusters, treated in a uniform way and with reliable errors. Early in this decade, this type of homogeneous study became possible for GCs (Paresce \\& De Marchi 2000), mostly thanks to the Hubble Space Telescope. In the meanwhile, high quality data have become available for YCs as well, mostly from wide field ground based surveys, thereby making this study possible on a global scale.  \\begin{center} \\begin{deluxetable*}{clllllrllcclc}[t] \\tablecaption{The sample of stellar clusters used in this study \\label{tab1}} \\tablehead{ \\colhead{ID} & \\colhead{Name} & \\colhead{$\\alpha$} & \\colhead{$\\beta$} & \\colhead{$m_c$} & \\colhead{$m$ range} & \\colhead{$[Fe/H]$} & \\colhead{$\\log M_{\\rm tot}$} & \\colhead{$c$} & \\colhead{$\\log t$} & \\colhead{$\\log t_{\\rm dis}$} & \\colhead{$\\tau_{\\rm dyn}$} & \\colhead{Ref} \\\\ \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{[M$_\\odot$]} & \\colhead{[M$_\\odot$]} & \\colhead{} & \\colhead{[M$_\\odot$]} & \\colhead{} & \\colhead{[yr]} & \\colhead{[yr]} & \\colhead{} & \\colhead{} } \\startdata 1 & $\\rho$\\,Oph   & $-1.9 \\pm 0.2$ & $1.9 \\pm 0.2$ & $0.17 \\pm 0.03$ & $0.02$ -- $7$ & $0.08$ & $3.3$ &   & $ 5.7$ & $8.7$&$0.001$ & a \\\\ 2 & Orion N. C.   & $-1.9 \\pm 0.2$ & $2.0 \\pm 0.3$ & $0.19 \\pm 0.05$ & $0.02$ -- $50$ & $-0.01$ & $1.8$ & $0.3$ & $ 6.0$ & $8.7$ &$0.002$ & b, c \\\\ 3 & Taurus        & $-2.3 \\pm 0.2$ & $1.8 \\pm 0.2$ & $0.51 \\pm 0.05$ & $0.03$ -- $3.5$& $-0.03$ & $1.7$  &       & $ 6.0$ & $6.3$ & $0.5$ & d \\\\ 4 & IC\\,348       & $-1.9 \\pm 0.1$ & $3.3 \\pm 0.2$ & $0.12 \\pm 0.05$ & $0.03$ -- $2.2$& $-0.01$ & $1.8$ & $0.3$ & $ 6.3$ & $8.7$ & $0.004$ & e \\\\ 5 & $\\sigma$\\,Ori & $-2.0 \\pm 0.1$ & $3.2 \\pm 0.2$ & $0.14 \\pm 0.02$ & $0.02$ -- $2.4$& $-0.01$ & $2.3$ & $0.6$ & $ 6.5$ & $8.7$ & $0.006$ & f \\\\ 6 & $\\lambda$\\,Ori& $-1.9 \\pm 0.1$ & $2.3 \\pm 0.2$ & $0.12 \\pm 0.02$ & $0.04$ -- $2.7$& $-0.01$ & $2.8$ & $0.6$ & $ 6.7$ & $8.7$ & $0.01$ & g \\\\ 7 & Cha\\,I & $-1.9 \\pm 0.2$ & $2.3 \\pm 0.3$ & $0.15 \\pm 0.03$ & $0.01$ -- $3.5$& $-0.11$ & $2.6$ &       & $ 6.7$ & $8.7$ &$0.01$ & h \\\\ 8 & IC\\,2391   & $-2.1 \\pm 0.1$ & $2.4 \\pm 0.3$ & $0.14 \\pm 0.02$ & $0.05$ -- $0.5$& $-0.09$ & $2.2$ & $0.6$ & $ 7.7$ & $8.8$ & $0.06$ & i \\\\ 9 & Blanco\\,1  & $-1.7 \\pm 0.1$ & $1.8 \\pm 0.2$ & $0.22 \\pm 0.03$ & $0.05$ -- $1.75$& $0.14$ & $3.5$ & $0.6$ & $ 8.1$ & $9.1$ & $0.09$ & j \\\\ 10 & Pleiades  & $-2.2 \\pm 0.1$ & $2.3 \\pm 0.2$ & $0.27 \\pm 0.03$ & $0.04$ -- $9$& $0.03$ & $2.9$ & $0.6$ & $ 8.1$ & $9.3$ & $0.06$ & k \\\\ 11 & M\\,35     & $-1.7 \\pm 0.1$ & $2.4 \\pm 0.2$ & $0.30 \\pm 0.03$ & $0.09$ -- $1.4$& $-0.21$ & $3.2$ & $0.7$ & $ 8.2$ & $9.4$ & $0.08$ & i \\\\ 12 & Coma Ber  & $-1.3 \\pm 0.4$ & $1.5 \\pm 0.3$ & $0.45 \\pm 0.10$ & $0.12$ -- $1.1$& $-0.08$ & $2.4$ & $0.4$ & $ 8.6$ & $8.9$ & $0.48$ & b \\\\ 13 & Praesepe  & $-2.0 \\pm 0.2$ & $4.0 \\pm 0.5$ & $0.20 \\pm 0.02$ & $0.07$ -- $1.5$& $0.00$ & $3.3$ & $0.7$ & $ 8.8$ & $9.5$ & $0.18$ & l, m\\\\ 14 & Hyades    & $-2.1 \\pm 0.1$ & $2.8 \\pm 0.2$ & $0.45 \\pm 0.05$ & $0.07$ -- $2.7$& $0.14$ & $2.6$ & $0.4$ & $ 8.8$ & $9.2$ & $0.37$ & n \\\\[0.15cm] 15 & NGC\\,104  & $-2.0$ & $1.8 \\pm 0.15$ & $0.33 \\pm 0.02$ & $0.11$ -- $0.7$ & $-0.76$ & $5.84$ & $2.03$ & $10.11$ & $10.93$& $0.15$ & p \\\\ 16 & NGC\\,5139 & $-2.0$ & $2.3 \\pm 0.2$  & $0.33 \\pm 0.02$ & $0.13$ -- $0.7$ & $-1.62$ & $6.17$ & $1.61$ & $10.06$ & $10.75$ & $0.21$ & p \\\\ 17 & NGC\\,5272 & $-2.0$ & $2.4 \\pm 0.2$  & $0.33 \\pm 0.02$ & $0.16$ -- $0.7$ & $-1.57$ & $5.65$ & $1.84$ & $10.06$ & $10.92$ & $0.14$ & p \\\\ 18 & NGC\\,6121 & $-2.0$ & $2.8 \\pm 0.25$ & $0.35 \\pm 0.02$ & $0.09$ -- $0.5$ & $-1.20$ & $4.83$ & $1.59$ & $10.10$ & $10.16$ & $0.88$ & p \\\\ 19 & NGC\\,6254 & $-2.0$ & $1.9 \\pm 0.2$ & $0.33 \\pm 0.02$ & $0.13$ -- $0.6$ & $-1.52$ & $5.00$ & $1.40$ & $10.06$ & $10.39$ & $0.47$ & p \\\\ 20 & NGC\\,6341 & $-2.0$ & $2.0 \\pm 0.2$ & $0.30 \\pm 0.02$ & $0.11$ -- $0.7$ & $-2.28$ & $5.30$ & $1.81$ & $10.12$ & $10.40$ & $0.52$ & p \\\\ 21 & NGC\\,6397 & $-2.0$ & $2.3 \\pm 0.2$ & $0.33 \\pm 0.02$ & $0.09$ -- $0.55$ & $-1.95$ & $4.65$ & $2.50$ & $10.10$ & $10.27$ & $0.67$ & p \\\\ 22 & NGC\\,6656 & $-2.0$ & $2.3 \\pm 0.3$ & $0.31 \\pm 0.03$ & $0.10$ -- $0.6$ & $-1.64$ & $5.46$ & $1.31$ & $10.10$ & $10.62$ & $0.31$ & p \\\\ 23 & NGC\\,6752 & $-2.0$ & $2.6 \\pm 0.1$ & $0.34 \\pm 0.02$ & $0.10$ -- $0.6$ & $-1.56$ & $5.15$ & $2.50$ & $10.07$ & $10.50$ & $0.37$ & p \\\\ 24 & NGC\\,6809 & $-2.0$ & $2.3 \\pm 0.3$ & $0.30 \\pm 0.02$ & $0.11$ -- $0.6$ & $-1.81$ & $5.04$ & $0.76$ & $10.09$ & $10.32$ & $0.58$ & p \\\\ 25 & NGC\\,7078 & $-2.0$ & $2.7 \\pm 0.3$ & $0.23 \\pm 0.02$ & $0.15$ -- $0.7$ & $-2.26$ & $5.75$ & $2.50$ & $10.07$ & $10.93$ & $0.14$ & p \\\\ 26 & NGC\\,7099 & $-2.0$ & $1.6 \\pm 0.2$ & $0.25 \\pm 0.02$ & $0.15$ -- $0.7$ & $-2.12$ & $5.00$ & $2.50$ & $10.11$ & $10.39$ & $0.53$ & p \\\\[0.15cm] 27 & NGC\\,2298 & $-2.0$ & $3.0 \\pm 0.2$ & $0.54 \\pm 0.01$ & $0.20$ -- $0.8$  & $-1.85$ & $4.50$ & $1.28$ & $10.02$ & $10.25$ & $0.59$ & q \\\\ 28 & NGC\\,6218 & $-2.0$ & $3.2 \\pm 0.4$ & $0.62 \\pm 0.03$ & $0.30$ -- $0.8$  & $-1.48$ & $4.94$ & $1.39$ & $10.10$ & $10.33$ & $0.60$ & q \\\\ 29 & NGC\\,6712 & $-2.0$ & $3.2 \\pm 0.2$ & $0.80 \\pm 0.03$ & $0.30$ -- $0.8$  & $-1.01$ & $4.97$ & $0.90$ & $10.02$ & $10.22$ & $0.63$ & q \\\\ 30 & NGC\\,6838 & $-2.0$ & $2.6 \\pm 0.2$ & $0.60 \\pm 0.02$ & $0.30$ -- $0.8$  & $-0.73$ & $4.23$ & $1.15$ & $10.14$ & $10.21$ & $0.84$ & q\\\\ \\enddata \\tablecomments{Table columns are as follows: ID number; cluster name; best fitting value of $\\alpha$ and $1\\,\\sigma$ uncertainty; best fitting value of $\\beta$ and $1\\,\\sigma$ uncertainty; best fitting value of $m_c$ and $1\\,\\sigma$ uncertainty; mass range over which the TPL fit is performed; cluster metallicity; total cluster mass; central concentration, or the  logarithmic ratio of tidal radius and core radius; cluster age; estimated time to dissolution; dynamical time, or the ratio of the cluster age and  the time to dissolution; bibliographic reference for the MF data, as follows. For YCs:  (a) \\cite{luh99}; (b) \\cite{kra07}; (c) \\cite{sle04}; (d) \\cite{sal10}; (e) \\cite{mue03}; (f) \\cite{cab08};  (g) \\cite{bar04a}; (h) \\cite{luh07}; (i) \\cite{bar04b}; (j) \\cite{mor07}; (k) \\cite{mor03}; (l) \\cite{cha05}; (m) \\cite{wil95}; (n) \\cite{bou08}. For dense GCs: (p) \\cite{pdm00}. For loose GCs: (q) \\cite{dem07}. For YCs, the metallicity and age are from the bibliography in column Ref. and references therein, whereas total mass and concentration are from Piskunov et al. (2008), { although the value of M$_{\\rm tot}$ is from K\\\"uc\\\"uk \\& Akkaya (2010) for Taurus, from Hodapp et al. (2009) for $\\sigma$\\,Ori and from Dolan \\& Mathieu  (2002) for $\\lambda$\\,Ori.} For GCs, the metallicity, total mass and concentration are from the 2003 revision of the GC catalogue of Harris (1996), whereas the age is from Forbes \\& Bridges (2010). The values of $t_{\\rm dis}$ have been determined as explained in Section\\,5, whereas $\\tau_{\\rm dyn}$ is the dynamical age of each cluster, expressed as $t/t_{\\rm dis}$.} \\end{deluxetable*} \\end{center} ",
        "conclusions": "At least qualitatively, the trend seen in Figure\\,\\ref{fig2} is fully consistent with the onset of mass segregation via two-body relaxation in a tidal environment, causing the preferential loss of low-mass stars from the cluster and hence a drift of the characteristic mass towards higher values. Vesperini \\& Heggie (1997) showed from N-body models that stellar evaporation, integrated over the cluster's orbit and further enhanced by the presence of the Galactic tidal field, causes a flattening of the MF, i.e. a selective depletion at the low-mass end. Their models do not reveal a drift of the characteristic mass towards higher values, but their IMF in all cases is a pure power-law. A drift of $m_c$ is visible in the N-body models of Portegies Zwart et al. (2001), although it proceeds very slowly and only appears when the study of the MF is limited to the regions inside the half-mass radius.  {That mass segregation is present in both young and globular clusters has long been established observationally. Examples include many of the YCs in our sample, such as M\\,35 (e.g. Mathieu 1983; Barrado y Navascu\\'es et  al. 2001), Pleiades (e.g. Converse \\& Stahler 2008), Hyades (Reid 1993; Perryman et al. 1998), Blanco\\,1 (Moreaux et al. 2007). As for GCs, mass segregation has been observed in all those objects where MF measurements exist at various radial distances (see e.g. De Marchi et al. 2007 and references therein) and can be satisfactorily explained as being the result of the two-body relaxation process (Spitzer 1987). Furthermore, De Marchi et al. (2007) have also shown that dense and loose GCs have today systematically different GMFs and this is consistent with a much stronger selective loss of low-mass stars in the latter. Loose GCs have systematically larger $m_c$ than denser clusters with otherwise similar metallicity or space motion parameters. This strongly suggests that the value of $m_c$ must drift as a result of the cluster's dynamical evolution, which is quicker in loose systems, even though the analysis of the N-body models (e.g. Portegies Zwart et al. 2001) has so far failed  to detect this effect.} In fact, Kruijssen (2009) has recently developed a simple physical model for the evolution of the MF that builds on the pioneering work of H\\'enon (1969). His results (see Figure\\,14 in Kruijssen 2009) show a shift in the inflection point of the MF with evolutionary time, consistent with a drift in $m_c$ (D. Kruijssen, priv. comm.).  { A potentially very interesting consequence of the trend shown in Figure\\,\\ref{fig2}, if dynamical evolution is indeed at its origin, is that it would seem plausible that all GCs, having now $m_c \\simeq 0.33$\\,M$_\\odot$ if they are dense or $\\sim 0.65$\\,M$_\\odot$ if they are loose, were born with a stellar MF like the one measured today in the youngest Galactic clusters, or more precisely with $m_c \\simeq 0.15$\\,M$_\\odot$. Although not necessarily required by Figure\\,\\ref{fig2}, this scenario is consistent with it.  The careful reader will notice that this statement appears at odds with the conclusions of a previous work by our team. From a study of the GMFs of twelve halo GCs with widely different dynamical histories and noting that they show a very similar characteristic mass ($m_c=0.33 \\pm 0.03$\\,\\Msolar), Paresce \\& De Marchi (2000) concluded that this value of $m_c$ had to be a feature of the stellar IMF in GCs. However, it is now apparent that this conclusion was the result of a selection effect, due to the fact that all twelve GCs in the Paresce \\& De Marchi (2000) sample have high central concentration, since no reliable information on the low-mass end of the MF of loose GCs was available at that time. The analysis presented by De Marchi et al. (2007) suggests that dense GCs have lost less low-mass stars than looser clusters throughout their life, but this does not mean that they have lost none. Therefore, one can no longer conclude that the $m_c \\simeq 0.33$\\,M$_\\odot$ value observed in the present-day GMF of dense GCs reflects the characteristic mass of their IMF, since it could have been smaller. }  The possibility that GCs were born with a much smaller $m_c$ value than what we measure today in their MF { would also be consistent} with the absence of a turn-over in the MF of the Galactic bulge down to the observational limit at $\\sim 0.2$\\,M$_\\odot$ (Zoccali et al. 2000). Regardless as to whether the bulge is primarily the result of fast formation at early epochs (e.g. Ballero et al. 2007) or a collection of the stars lost from disrupted clusters throughout the life of the Galaxy (e.g. Gnedin \\& Ostriker 1997), its very location at the bottom of the Galaxy's potential well makes it very hard for low-mass stars to escape from it. Therefore, while the present MF of the Galactic disc could represent a complex average over time and space, that of the bulge should have not been altered by dynamical evolution and, unlike the case of the comparably old GCs, should still reflect the properties of the IMF. Forthcoming HST observations of the Galactic bulge (see e.g. Brown et al. 2009) will provide insights on its MF below $0.15$\\,M$_\\odot$, where we expect a turn-over.  {To be sure, N-body models of dense stellar systems able to deal with a large number of particles ($N > 10^6$) and incorporating a realistic treatment of stellar evolution and of the interaction with the Galactic tidal field (see Portegies Zwart et al. 2010) would be necessary to simulate the growth of the characteristic mass with time in GCs. Our analysis (see Section\\,5) suggests that young and globular clusters may share rather similar values of the $\\alpha$ and $\\beta$ parameters. Thus, if the picture sketched above is correct and, by strongly affecting the value of $m_c$, dynamical evolution act as the main source of variations in the present-day stellar MF of Galactic clusters, one might postulate that Population I and II stars should have similar IMFs. While not necessarily required by our analysis, this hypothesis is fully compatible with our results. Claims of the IMF universality that one finds in the literature (e.g. Gilmore 2001; Elmegreen et al. 2008) are based on the observation that stellar MFs in different environments have {\\em similar} $m_c$ values, within a factor of a few. Here we show that it is plausible that the original value of $m_c$ be actually {\\em the same} for Population I and II stars.  If this is true, one would have to consider the possibility that the characteristic mass in the IMF is not necessarily set by the Jeans mass scale in the forming cloud (e.g. Larson 1998; Larson 2005) and that the physical conditions of the environment do not play a significant role in the outcome of the star formation process, i.e. on the IMF. This, however, would not necessarily mean that the star formation process itself is the same: recent numerical simulations proposed by Goodwin \\& Kouwenhoven (2009) and Bate (2009) suggest that the IMF shape is largely insensitive to parameters such as the binary fraction and mass ratio in binaries. If this is the case, the role of the IMF as an effective tool to probe the physics of star formation may have to be critically reconsidered.}"
    },
    "1005/1005.4070_arXiv.txt": {
        "abstract": "{ We present evidence for the contribution of high-mass globular clusters to the stellar halo of the Galaxy. Using SDSS-II/SEGUE spectra of over 1900 G- and K-type halo giants, we identify for the first time a subset of stars with CN bandstrengths significantly larger, and CH bandstrengths lower, than the majority of halo field stars, at fixed temperature and metallicity. Since CN bandstrength inhomogeneity and the usual attendant abundance variations are presently understood as a result of star formation in globular clusters, we interpret this subset of halo giants as a result of globular cluster dissolution into the Galactic halo. We find that $2.5\\%$ of our sample is CN-strong, and can infer based on recent models of globular cluster evolution that the fraction of halo field stars initially formed within globular clusters may be as large as $50\\%$.} ",
        "introduction": "Hierarchical structure formation is presently the dominant explanation for galaxy formation, based on observed fluctuations in the cosmic microwave background \\citep{WMAP} and sophisticated numerical simulations of their evolution to the present day (e.g., Diemand et al. 2007\\nocite{DKM07}; Springel et al. 2008\\nocite{SWV08}). In this picture, galaxies like the Milky Way are formed through the coalescence of multiple low-mass galaxies which develop within a much larger dark matter halo. The initial disagreement between the calculated mass function of dark matter subhaloes in the simulations and the observed mass function of nearby dwarf galaxies (e.g., Klypin et al. 1999\\nocite{KKV99}; Moore et al. 1999\\nocite{MGG99}) is being addressed from several directions at the same time, from complex semianalytic simulations that include star formation and feedback processes, and calculate the chemodynamical evolution of Milky Way-like galaxies (e.g., Johnston et al. 2008\\nocite{JB08}; Tumlinson 2010\\nocite{T10} and references therein) to searches for extremely low-mass galaxies in the Local Group (e.g., Zucker et al. 2006a, 2006b\\nocite{ZB06a}\\nocite{ZB06b}; Belokurov et al. 2007\\nocite{BZE07}).   The stellar halo of the Milky Way is thought to have been constructed mostly through the early ($\\simeq 10$ Gyr ago) accretion of low-mass protogalaxies. The halo exhibits considerable substructure in density and in velocity (e.g. Bell et al. 2008\\nocite{BZB08}), and kinematically distinct streams presently observed in the halo (e.g., Majewski et al. 2003\\nocite{MSW03}; Duffau et al. 2006\\nocite{DZV06}; Mart{\\'{\\i}}nez-Delgado et al. 2007\\nocite{MD07}) are interpreted as remnants of more recent or ongoing merger activity. The ``ECHOS'' identified in \\citet{SRB09} are interpreted as older substructure that has lost some spatial coherence with time. Studies of the abundance distributions and star formation histories in nearby dwarf galaxies (e.g., Koch et al. 2007a, 2007b, 2008a, 2008b\\nocite{KG07}\\nocite{KW07}\\nocite{KG08}\\nocite{KM08}; Kirby et al. 2008\\nocite{KSG08}; Aoki et al. 2009\\nocite{AAS09}; Frebel et al. 2009\\nocite{FKS09}, 2010\\nocite{FSGW10}) have shown that there is a reasonable concordance between the properties of the present-day Milky Way halo (as characterized by, e.g., Schoerck et al. 2009\\nocite{SCC09}) and the dwarf galaxies that would have been available as stellar contributors early in Galactic history (e.g., Font et al. 2006\\nocite{FJB06}; Carollo et al. 2007\\nocite{CB07}).  However, the ongoing dissolution of globular clusters such as Palomar 5 (e.g., Odenkirchen et al. 2001, 2002, 2003\\nocite{OGR01}\\nocite{OGD02}\\nocite{OGD03}; Rockosi et al. 2002\\nocite{ROG02}; Grillmair \\& Dionatos 2006\\nocite{GD06}) and NGC 5466 \\citep{BEI06} implies that some fraction of halo stars are initially formed in globular clusters. This dissolution is driven by internal 2-body relaxation, stellar evolution processes and tidal interactions with the Galactic potential, and can cause significant mass loss over the lifetime of typical halo globular clusters (e.g., Gnedin \\& Ostriker 1997\\nocite{GO97}).  One particular model of globular cluster formation, described in \\citet{DVD08}, posits that all old halo clusters surviving to the present day have lost at least 90\\% of their initial mass. Early in the development of the cluster, winds from AGB stars collected in the cluster center and formed a second generation of low-mass stars. Type Ia supernovae then caused the cluster to expand, and stars at large radii, mostly members of the first generation, were lost. In the model, the end result, $\\simeq 10$ Gyr later, is a $10^{5} - 10^{6} M_{\\sun}$ cluster with two stellar populations that differ slightly in age and abundance pattern.  This two-generation model was developed specifically to explain the anomalous light-element abundance patterns observed in globular cluster stars, specifically the presence in every old globular cluster in the Milky Way of a subpopulation with typical Population II abundances, and a second subgroup with the same metallicity but enhanced N, Na, and Mg along with depleted C, O, and Al. This abundance bimodality has been studied extensively in globular clusters (Langer et al. 1992\\nocite{LSK92}, Kraft 1994\\nocite{K94}, and Gratton et al. 2004\\nocite{GSC04} all provide thorough reviews of the topic), and explanations for the second abundance subgroup have varied from pollution \\citep{CC98} to internal mixing \\citep{L85} to enrichment of star-forming gas by moderate- to high-mass stars (e.g., Cottrell \\& DaCosta 1981\\nocite{CD81}; Yong et al. 2008\\nocite{Y08}) or high-mass binaries \\citep{DMP09}. Surface pollution of already-formed stars would result in abundance anomalies that are erased at first dredge-up, while current models of deep mixing (e.g., Charbonnel \\& Zahn 2007\\nocite{CZ07}) indicate that it only begins to operate after first dredge-up. Both of these processes have been effectively ruled out as explanations for abundance bimodality by the presence of abundance variations at all evolutionary phases in globular clusters (e.g., Briley et al. 2002\\nocite{BCS02}; Harbeck et al. 2003\\nocite{HSG03a}). In the primordial enrichment scenario, there is ongoing discussion over the exact source of enriching material. In some models the source is moderately high-mass ($\\sim 4-5M_{\\sun}$) AGB stars (e.g., Parmentier et al. 1999\\nocite{PJM99}), while \\citet{DCM07} claim that rotating high-mass ($M \\ga 10 M_{\\sun}$) stars are a better source for processed material because of their very short lifetimes and \\citet{DMP09} prefer high-mass binaries because of their potentially strong mass loss and low wind velocities.  Stars with these light-element abundance anomalies are readily identified through strong UV/blue CN molecular absorption and relatively weak absorption in the CH G band, and are hereafter called ``CN-strong stars'', with the understanding that the full abundance pattern from C through Al necessarily follows the CN variation. They are not observed to exist in open clusters (e.g., Smith \\& Norris 1984\\nocite{SN84}; Jacobson et al. 2008\\nocite{JFP08}; Martell \\& Smith 2009\\nocite{MS09}) or the halo field \\citep{GSC04}. This feedback process apparently only occurs in the high-density environment of globular clusters, and as a result the characteristic sawtooth abundance pattern from carbon through aluminium can be used as a marker of globular cluster origin.  Given the contributions globular clusters are presently making to the halo field, and the significant mass loss predicted theoretically over the lifetime of the Galactic globular cluster system (e.g., Baumgardt et al. 2008\\nocite{BKP08}), it is intriguing that no CN-strong stars have to date been observed in the halo. We interpret this as a qualitative sign that the contributions to the halo field of globular clusters as we know them today are relatively minor, as is also suggested in \\citet{Y08}.  To test this interpretation, we searched for CN-strong halo giants in the Sloan Extension for Galactic Understanding and Exploration (SEGUE) survey \\citep{Y09}. The SEGUE survey is a spectroscopic extension of imaging taken during the Sloan Digital Sky Survey \\citep{SDSS00}, with targets selected to address questions of halo substructure and Galactic formation history. Data were taken from 2005 August through 2008 July using a 640-fiber multiobject spectrograph and the same telescope at Apache Point Observatory that was used for SDSS imaging. The first portion of the SEGUE data was made publically available in 2008 as part of SDSS Data Release 7 (DR7), including roughly 240,000 spectra in 200 ``pencil beam'' lines of sight containing stars chosen for specific purposes (i.e., fields in the Sagittarius stream, M dwarfs to study extremely local kinematic substructure, G and K giants to study the distant halo). In addition to flux-calibrated spectra, DR7 also offers the products of the ``SEGUE Stellar Parameters Pipeline'' (SSPP), which include derived stellar parameters like effective temperature, [Fe/H] metallicity, and radial velocity, determined automatically through template matching, $\\chi^{2}$ minimization, cross-correlation or grid-matching methods, as appropriate. Lee et al. (2008a\\nocite{L08a}; 2008b\\nocite{L08b}) and \\citet{AP08} give thorough explanations of the SSPP pipeline and process. ",
        "conclusions": "Although light-element abundance variations have not been observed before in the halo, our identification of these CN-strong halo field stars is not wholly unexpected. There are several well-understood mechanisms for globular cluster mass loss, and theoretical studies of globular cluster formation and evolution (e.g., D'Ercole et al. 2008\\nocite{DVD08}; Baumgardt et al. 2008\\nocite{BKP08}) predict significant mass loss in individual clusters as well as a dramatic reshaping of the cluster mass function with time. The data set we analyzed, selected from the SEGUE survey, is not representative of the full halo, in mass, evolutionary phase, or metallicity. Distances to the candidate CN-strong stars range from 4 to nearly 40 kpc, with $93\\%$ found within 20 kpc of the Sun. However, our result for red giants is generalizable to all halo stars, since abundance bimodality is observed to exist at all masses and evolutionary phases in globular clusters. More fundamentally, all of the cuts we made in selecting the final data set are blind to CN and CH bandstrengths and light-element abundances, and all nitrogen-enhanced giants with metallicities above [Fe/H]$=-1.8$ ought to show clearly strong CN bands. Since approximately $2.5\\%$ of our halo red giants exhibit strong CN bands and weak CH bands, we expect that the same fraction of the entire halo will contain the same abundance enhancements and depletions. This prediction can be confirmed by observations of dwarfs in the halo field, because main sequence stars in globular clusters show the same abundance division as giants (e.g., Briley et al. 2002\\nocite{BCS02}).  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{f9.eps}} \\caption[Generalized histograms, all metallicities]{ Generalized histograms of $\\delta S(3839)$ for CN-weak and CN-strong stars in each [Fe/H] bin. As in Fig. \\ref{ff6}, maximum metallicity for each panel is given in the upper left corner. As in Fig. \\ref{ff5}, CN-strong histograms were amplified for clarity. The multiplicative factor for each panel is given in the upper right corner. } \\label{ff9} \\end{figure}  In order to convert $f^{\\mathrm{p}}_{\\mathrm{h}}$, the present-day fraction of CN-strong halo stars, into $f^{\\mathrm{gc}}_{\\mathrm{h}}$, the fraction of globular cluster-originating stars in the halo field, we must consider what fraction CN-strong stars comprise of the stars originally formed in globular clusters. In the two-generation scenario of \\citet{DVD08}, roughly $90\\%$ of stars originally formed in a globular cluster, consisting entirely of first-generation stars (with halo-like chemistry), are lost between the epoch of cluster formation and the present day. This means that $f^{\\mathrm{m}}_{\\mathrm{gc}}$, the fraction of stars that remain as members of the globular cluster they were formed in, is around $0.1$. Since $f^\\mathrm{p}_\\mathrm{gc}$, the fraction of present-day globular clusters stars that are CN-strong, is around $0.5$ (e.g., Kraft 1994\\nocite{K94}), we can calculate that $f^{\\mathrm{gc}}_{\\mathrm{h}}=\\frac{f^{\\mathrm{p}}_{\\mathrm{h}}}{f^{\\mathrm{m}}_{\\mathrm{gc}}\\times f^{\\mathrm{p}}_{\\mathrm{gc}}}$. Since $f^{\\mathrm{p}}_{\\mathrm{h}}=0.025$ in the present study, this suggests that $f^{\\mathrm{gc}}_{\\mathrm{h}}=\\frac{0.025}{0.1\\times 0.5}=0.5$, and that a remarkable $50\\%$ of the halo field originally formed in the massive star clusters that were progenitors of the present-day globular cluster population, with a further unknown contribution of CN-weak stars made by globular clusters that did not survive to the present day and were not massive enough to self-enrich.  While some numerical studies of galaxy formation (e.g., Boley et al. 2009\\nocite{BLR09}) have suggested that the halo could be constructed entirely from disrupted globular clusters, there is not presently a strong consensus on the role of globular clusters in cosmological-scale galaxy formation. Precise numerical study of the dynamical evolution of globular clusters is very complicated: the number of particles is large enough, and the relevant timescales short enough, that highly accurate simulations are very time-consuming. However, the development of semianalytic prescriptions for the mass evolution of globular clusters would allow single-halo-scale simulations like those of \\citet{JB08} to include them as a source of halo stars, and to predict what fraction of the halo field ought to originate in globular clusters."
    },
    "1005/1005.4888_arXiv.txt": {
        "abstract": "The daily variation of the solar photocenter over some 11 years is derived from the Mount Wilson data reprocessed by Ulrich et al. 2010 to closely match the surface distribution of solar irradiance. The standard deviations of astrometric jitter are 0.52 $\\mu$AU and 0.39 $\\mu$AU in the equatorial and the axial dimensions, respectively. The overall dispersion is strongly correlated with the solar cycle, reaching $0.91~\\mu$AU at the maximum activity in 2000. The largest short-term deviations from the running average (up to 2.6 $\\mu$AU) occur when a group of large spots happen to lie on one side with respect to the center of the disk. The amplitude spectrum of the photocenter variations never exceeds $0.033$ $\\mu$AU for the range of periods $0.6$--$1.4$ yr, corresponding to the orbital periods of planets in the habitable zone. Astrometric detection of Earth-like planets around stars as quiet as the Sun is not affected by star spot noise, but the prospects for more active stars may be limited to giant planets. ",
        "introduction": "\\label{firstpage}  The prospects of finding habitable planets orbiting nearby solar-type stars are to a large degree associated with the ultra-precise astrometric instruments under development or construction, such as the SIM Observatory \\citep{sha,unw,catpa} and Gaia \\citep{cas}. The Earth orbiting the Sun produces an observable astrometric wobble of 3 $\\mu$AU (micro-AU) and a radial velocity variation of 0.089 \\ms, if seen equator-on (inclination $i=90\\degr$). A pole-on configuration ($i=0\\degr$ or $180\\degr$) is optimal for astrometry, because the reflex motion signal is present in both dimensions of the sky projection, whereas the radial velocity amplitude, which is proportional to $\\sin i$, drops to zero. For a high signal-to-noise (S/N) ratio detection, the total error budget of the prospective exoplanet detection techniques should be well below these values. The astrophysical jitter caused by the rotation of magnetic features on the surface (spots, faculae) potentially becomes a significant part of the total observational error in this domain of accuracy, even for fairly common inactive dwarfs. The recently completed double blind test of planet-detection capabilities with the SIM Observatory produced encouraging results even for complex multiple systems \\citep{tra}, but magnetic jitter was not taken into account in those simulations. Semi-analytical considerations in \\citep{mak09} supported by Monte-Carlo simulations and indirect observational evidence determined that the astrometric method is much less sensitive to the effects of magnetic features, at least by an order of magnitude for solar-type stars, than the Doppler shift technique. Similar conclusions were drawn by \\citet{cat} for the Sun, based on a sophisticated model of sunspot activity and extensive numerical simulation.  The variable distribution of surface brightness is the cause of both the astrometric jitter and the total flux variation. Therefore, the former can be estimated from the latter, given a model describing the number of spots (or bright features), their size and lifetimes. The dispersion of the total solar irradiance (TSI) observed for a few decades with a number of satellites \\citep{fro} sets the standard of quiet stars ($4\\cdot 10^{-4}$, or 400 ppm in normalized flux). The light curves from Kepler for some 120000 stars revealed that at least half of all solar-type dwarfs in the field are at least as quiet as the Sun \\citep{bas}, displaying lower levels of photometric dispersion. By extrapolation, about half of nearby dwarfs should be as amenable to exoplanet detection as the Sun. The aim of this paper is to determine the solar astrometric jitter directly from observations. The results can be used to verify and correct the existing models of spot activity for other stars. ",
        "conclusions": "We determined that the solar jitter caused by magnetic features is small and should not preclude astrometrists from detecting habitable planets as small as the Earth. A few additional considerations should be made when extrapolating this result to other stars or to actual astrometric observations.  \\begin{figure}[htbp] \\plotone{SunHZdX_power.eps} \\caption{Amplitude spectrum of solar astrometric jitter in the equatorial dimension for the range of periods corresponding to the habitable zone.} \\label{pow.fig} \\end{figure}  The ultra-precise photometric measurements with Kepler confirmed that about half of all solar-type dwarfs are as quiet as the Sun, or more \\citep{bas}. The stars that are more active than the Sun have much larger filling factors, and their magnetic features may be long-lived (up to a few weeks). Both photometric and astrometric variability is proportional to the average total area $A_s$ occupied by spots \\citep{mak09}. In the model currently adopted for planet detection simulations with SIM (to be published elsewhere), the total area is related to the index of chromospheric activity $\\log R'_{\\rm HK}$ through \\eb \\log A_s=12.468+1.75\\log R'_{\\rm HK}, \\ee which yields 6137 micro-stellar-hemispheres (MSH) for the Sun. As roughly half of nearby field stars are more chromospherically active than the Sun \\citep{gra}, they are expected to have larger spot groups and, therefore, jitter amplitudes. At $\\log R'_{\\rm HK}=-4.6$, the expected area and the jitter are 5 times the solar values, and at $\\log R'_{\\rm HK}=-4.2$, 25 times. This model gives fairly accurate predictions for the active star $\\kappa^1$ Ceti, carefully studied with MOST \\citep{wal}. Furthermore, an additional factor of 2--6 increase may ensue from the limited duration of an astrometric mission. Both SIM and Gaia will nominally operate for 5 years. This is considerably shorter than the solar cycle. If a solar analog is observed at the height of its activity, the level of magnetic jitter may be higher than the multi-year average. Thus, magnetic jitter can indeed pose the natural limit to the sensitivity of the astrometric method at stars significantly more active than the Sun.  The amplitude spectrum periodogram shown in Fig.~\\ref{pow.fig} is derived from a high-frequency observational cadence of 2881 data points over 11 years. We expect only about 200 2D measurements to be taken for each target star with SIM Lite during its 5-year nominal mission. If the astrometric jitter is completely random and uncorrelated, the variance of the amplitude of a random harmonic in the spectrum is inversely proportional to the number of data points. In that case, the amplitude spectrum will rise by a factor of $\\sqrt{2881/200}=3.8$ for SIM. On the other hand, if the magnetic activity underlying the astrometric jitter is systematic and deterministic, the corresponding parts of the spectrum will stay as low is in Fig.~\\ref{pow.fig}. Such systematic variations may occur, for example, if sunspots and plages are not randomly distributed on the solar surface, but tend to appear in confined areas called ``active longitudes\" in the literature. Currently, there seems to be no strong evidence for the existence of such long-term structures. A possible north-south asymmetry in the distribution of magnetic features evolving with the solar cycle is another interesting route of investigation where the reconstructed TSI images can be used.  The results derived in this paper apply to solar-type stars seen equator-on (inclination $\\simeq 90\\degr$). A pole-on configuration is optimal for astrometric detection, because both dimensions are equally engaged, but it is statistically less probable. Even at moderate inclinations to the line of sight, the distribution of magnetic perturbations in $y$-coordinate may be different from that we find for the Sun, if the magnetic features are confined to the equatorial zone. The features would normally be seen closer to the side of the limb opposite to the visible pole, producing a skewed distribution of $\\Delta y$ depending on whether the spots or bright areas are the dominating contributors. However, giant near-polar spots appear to be common on very active, fast-rotating stars."
    },
    "1005/1005.1180_arXiv.txt": {
        "abstract": "We use the UKIDSS Ultra-Deep Survey to trace the evolution of galaxy clustering to $z=3$. Using photometric redshifts derived from data covering the wavelength range $0.3 - 4.5 {\\rm \\mu m}$ we examine this clustering as a function of absolute K-band luminosity, colour and star-formation rate. Comparing the deprojected clustering amplitudes, we find that red galaxies are more strongly clustered than blue galaxies out to at least $z=1.5$, irrespective of rest-frame K-band luminosity. We then construct passive and star-forming samples based on stellar age, colour and star-formation histories calculated from the best fitting templates. The clustering strength of star-forming galaxies declines steadily from $r_0\\simeq7h^{-1}$Mpc at $z\\simeq2$ to $r_0\\simeq3h^{-1}$Mpc at $z\\simeq0$, while passive galaxies have clustering strengths up to a factor of two higher. Within the passive and star-forming subsamples, however, we find very little dependence of galaxy clustering on K-band luminosity. Galaxy `passivity' appears to be the strongest indicator of clustering strength. We compare these clustering measurements with those predicted for dark matter halos and conclude that passive galaxies typically reside in halos of mass $M\\ge10^{13}M_{\\odot}$ while luminous star-forming galaxies occupy halos an order of magnitude less massive over the range $0.5<z<1.5$. The decline in the clustering strength of star-forming galaxies with decreasing redshift indicates a decline in the hosting halo mass for galaxies of a given luminosity. We find evidence for convergence of clustering in star-forming and passive galaxies around $z\\sim2$, which is consistent with this being the epoch at which the red sequence of galaxies becomes distinct. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.3699_arXiv.txt": {
        "abstract": "s{The existence of cosmic rays and weak magnetic fields in the intracluster volume has been well proven by deep radio observations of galaxy clusters. However a detailed physical characterization of the non-thermal component of large scale-structures, relevant for high-precision cosmology, is still missing. I will show the importance of combining numerical and theoretical works with cluster observations by a new-generation of radio, Gamma- and X-ray instruments.  } ",
        "introduction": "Deep radio observations of the sky have revealed the presence of extended ($\\sim$ 1 Mpc) radio sources in about 50 merging galaxy clusters (see \\cite{Feretti08,2008SSRv..134...93F} and references therein). This diffuse radio emission is not related to unresolved radio-galaxies, but rather to the presence of relativistic particles ($\\gamma >>$1000) and magnetic fields of the order of $\\mu$Gauss in the intracluster volume. The physical mechanisms responsible for the origin of this non-thermal intracluster component are matter of debate (e.g. \\cite{1977ApJ...212....1J,Govoni04}), as well as the effects of intracluster cosmic rays (CRs) and magnetic fields on the thermodynamical evolution and mass estimate of galaxy clusters (e.g. \\cite{2009arXiv0909.0270S,2010A&A...510A..76L}).  A deep understanding of the evolutionary physics of {\\em all} the different cluster components (dark matter, galaxies, thermal and non-thermal intracluster medium -- ICM) and of their mutual interactions is indeed essential for high-precision cosmology with galaxy clusters \\cite{2005bmri.conf...77A}.  In the following, I will give an overview of our current knowledge of the non-thermal component of galaxy clusters. I will also stress the importance of a new generation of multi-wavelength telescopes -- such as the {\\it Low Frequency Array} ({\\it LOFAR}), and the Gamma- and hard X-ray (HXR) satellites {\\it Fermi} and {\\it NuSTAR} -- for a deep understanding of the non-thermal cluster physics. The $\\Lambda$CDM model with H$_0$=70 km ${\\rm s}^{-1} {\\rm Mpc}^{-1}$, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$ has been adopted. ",
        "conclusions": ""
    },
    "1005/1005.4907_arXiv.txt": {
        "abstract": "We investigate the role of host galaxy classification and black hole mass ($M_{BH}$) in a heterogeneous sample of 276 mostly nearby ($z<0.1$) X-ray and IR selected AGN. Around 90$\\%$ of Seyfert 1 AGN in bulge-dominated host galaxies (without disk contamination) span a very narrow range in the observed 12$\\micron$ to 2-10keV luminosity ratio ($1<R_{IR/X}<7$). This narrow dispersion incorporates \\emph{all} possible variations among AGN central engines, including accretion mechanism and efficiency, disk opening angle, orientation to sightline, covering fraction of absorbing material, patchiness of X-ray corona and measured variability. As a result, all models of X-ray and IR production in AGN are very strongly constrained. Among Seyfert 1 AGN, median X-ray and IR luminosities increase with black hole mass at $>99\\%$ confidence. Using ring morphology of the host galaxy as a proxy for lack of tidal interaction, we find that AGN luminosity in host galaxies within 70Mpc is independent of host galaxy interaction for $\\sim$ Gyrs, suggesting that the timescale of AGN activity due to secular evolution is much shorter than that due to tidal interactions. We find that LINER hosts have lower 12$\\micron$ luminosity than the median $12\\micron$ luminosity of normal disk- and bulge-dominated galaxies which may represent observational evidence for past epochs of feedback that supressed star formation in LINER host galaxies. We propose that nuclear ULXs may account for the X-ray emission from LINER 2s without flat-spectrum, compact radio cores. We confirmed the robustness of our results in X-rays by comparing them with the 14-195keV 22-month BAT survey of AGN, which is all-sky and unbiased by photoelectric absorption. ",
        "introduction": "\\label{sec:intro} Galactic nuclei generally considered to be 'active' have bolometric luminosities of $\\sim 10^{-2}$ to $\\sim 10^{4}$ times their host galaxy luminosity and are powered by accretion onto a supermassive black hole. The fundamental parameters that determine accretion onto the central black hole should be: (1) the black hole mass, (2) the amount of gas and dust in the galactic nucleus, (3) a mechanism to drive material in the galactic nucleus onto the black hole and (4) outflows as a result of accretion, which may affect (2) and (3) as feedback. While most galactic nuclei in the local Universe are believed to host very large mass black holes, only $\\sim 0.01-1\\%$ of these nuclei are highly luminous \\citep[e.g.][]{b49,b13}. By contrast, low luminosity nuclei are far more common, accounting for $\\sim 1/3$ of all galactic nuclei \\citep{b95}. Therefore some combination of the parameters determining accretion conspires to keep very luminous galactic nuclei relatively rare, while also providing a very wide range of observed activity at lower luminosities.  Black hole mass is the simplest parameter involved in accretion. The greatest activity (or highest bolometric luminosity) in a galactic nucleus occurs when a black hole accretes mass quickly. Black hole mass correlates with several properties of the central regions of the host galaxy, including bulge stellar velocity dispersion ($\\sigma_{\\ast}$), bulge mass and bulge luminosity (see e.g. \\citet{b39} and references therein). This implies that the largest mass black holes live in galaxies with large bulges. If most galactic bulges grew early in the Universe \\citep[e.g.][]{b41}, then black holes grew fastest at that time and so activity in galactic nuclei may have been greatest at high redshift. However, periods of intense activity in a galactic nucleus could also arise at lower redshift, due to mergers between gas-rich galaxies \\citep[e.g.][]{b43} or when a large reservoir of cold gas builds up in the galactic disk \\citep{b24,b40}.  The raw material for accretion onto the black hole is gas and dust in the galactic nucleus. How much gas and dust there is in the galactic nucleus depends on the local rate of formation of massive stars \\citep{b25} and on mechanisms driving gas and dust into the nucleus from elsewhere in the galaxy. Mechanisms driving material into the galactic nucleus from the outside could be internal or external to the galaxy. Internal mechanisms involve bars \\citep[e.g.][]{b33} or more generally, internal (secular) disk driven evolution \\citep{b24}; external mechanisms involve tidal disruptions or mergers \\citep[e.g.][]{b41} or nuclear bombardment \\citep{b60}. Not all of the material in the galactic nucleus needs to come from the rest of the galaxy. Material that gains angular momentum close to the accreting supermassive black hole can flow outwards to be recycled in the surrounding galactic nucleus \\citep[e.g.][]{b27,b1,b44,b8}. Indeed if the feedback from the accreting black hole is powerful enough it could disrupt star formation in the galactic nucleus and beyond \\citep{b1}. Outflows from the accreting black hole into the surrounding galactic nucleus could end up terminating inflows onto the black-hole itself, which leads to a picture of galactic nucleus activity as self-regulating \\citep[e.g.][]{b38}.  Isolating the fundamental parameters that determine the mass accretion rate onto supermassive black holes is a major observational problem. Many active galactic nuclei (AGN) are shrouded by obscuring material \\citep{b2}, although the obscuration can be complicated \\citep[e.g.][]{b4,b3}. AGN are mostly distant enough that broad-band observations include host galaxy luminosity contributions (e.g. from hot diffuse gas, X-ray binaries and ultra-luminous X-ray sources in the X-ray band alone). One approach to solving this difficult observational problem is to compare broadband luminosities in AGN with fundamental accretion parameters, such as black hole mass, or simple observables such as host galaxy classification, which may be related to fundamental accretion parameters.  In this work, we investigate the connection between black hole mass, host galaxy classification and the observed IR and X-ray luminosities of a heterogeneous sample of 276 AGN (mostly from \\citep{b99}). In section~\\ref{sec:sample} we investigate the connection between AGN luminosity, black hole mass and host galaxy classification. In section~\\ref{sec:hockey} we discuss the AGN luminosity distribution and in section~\\ref{sec:group2} we discuss the highly heterogeneous, non-Seyfert 1 AGN in our sample. In section~\\ref{sec:rings}, we discuss the importance of ringed morphology in host galaxies and the implications for AGN activity. We discuss issues of bias and completeness in our sample as well as the reliability of our conclusions in section~\\ref{sec:bias} and in section~\\ref{sec:conclusions} we summarize our conclusions. ",
        "conclusions": "\\label{sec:conclusions} We investigated the role of black hole mass ($M_{BH}$) and host galaxy morphology in a survey of the observed 2-10keV X-ray and 12$\\micron$ IR luminosities of a heterogeneous sample of 276 mostly nearby AGN. We find that:\\\\  (a) As black hole mass increases, the average observed IR and X-ray luminosity of the homogeneous Group 1 (Seyfert 1.X) AGN increases, maintaining a ratio in the range $R_{IR/X}=$[1,30] over $\\sim 3-4$ orders of magnitude in mass at a confidence level of $>90\\%$. The luminosities for Group 1 AGN remain in the range $\\sim 10^{-3}-10^{-1}$ of Eddington for each decade in mass. By contrast, among the heterogeneous, lower luminosity Group 2 AGN, the ratio $R_{IR/X}$ \\emph{decreases} as black hole mass increases. This is mostly due to a combination of host galaxy contamination and photoelectric absorption.  (b) There is an average increase in IR luminosity of a factor of $\\sim 3$ among Group 1 AGN in disk-dominated hosts versus bulge-dominated hosts at a confidence level $>99\\%$. Presumably the increase is due to contamination from the unresolved galactic disk \\citep{b12}. Therefore a better measure of the underlying central engine is the range of $R_{IR/X}$ in bulge-dominated Group 1 AGN, which is $1<R_{IR/X}<7$ at $\\sim$90$\\%$ confidence. This dispersion incorporates \\emph{all} possible variations among these Group 1 AGN, ranging from accretion mechanism and efficiency, to the opening angle of the disk and orientation to sightline, to covering fraction of absorbing material, to patchiness of the X-ray corona and even including variability. Group 1 AGN are quite remarkably identical. This strongly constrains all models of X-ray and IR production in AGN.  (c) Among the lower luminosity Group 2 AGN, those in bulge-dominated hosts tend to have substantially lower $L_{IR}$ and/or $R_{IR/X}$ compared to those in disk-dominated galaxies. Among LINERs we find much lower $L_{IR}$ than in normal disk-dominated or bulge-dominated galaxies representing possible evidence for past epochs of AGN feedback in both bulge and disk-dominated hosts.  (d) Within 70Mpc, AGN in host galaxies with outer Linblad resonance rings (R,R') are statistically indistinguishable in $L_{X},L_{IR}$ and $M_{BH}$ from AGN in host galaxies without rings. Since these rings are hard to form and easily destroyed by perturbations \\citep{b31,b80,b73}, the ringed host galaxies are unlikely to have tidally interacted with companions for a long time.  Since AGN activity typically lasts a few $\\times 10$Myrs, the most likely explanation for this is that the timescale for activity driven by processes internal to the galaxy ($t_{int}$) is much smaller than the timescale for activity driven by processes external to the galaxy ($t_{ext}$), or $t_{int}\\ll t_{ext}$.  (e) \\emph{ULXs in galactic nuclei could account for the observed X-ray luminosity in certain LINERs}. We propose that the X-ray emission from LINER 2 nuclei without flat-spectrum, compact radio cores \\citep{b5} originates in ULXs (most likely large stellar mass black holes accreting at near Eddington luminosity). This hypothesis is testable by searching for a combination of short timescale X-ray variability, large fractional variability (the excess variance) and spectral changes in LINER 2 nuclei without these radio properties. A large estimate for the central mass in this case may simply indicate the presence of a quiescent supermassive black hole. Our proposal may be consistent with the 'retired galaxies' of \\citet{b21} and could be tested in that context in a study of the optical spectra of host galaxies of the 'radio quiet' LINER 2s.  (g) We tested results (a-d) above by comparing the 14-195keV luminosity and the 12$\\micron$ IRAS luminosity of AGN in the BAT 22-month sample \\citep{b57}. We conclude that photoelectric absorption is not important in the 2-10keV band for Group 1 AGN, but it is important for many (though not all) of the Group 2 AGN and accounts for some of the 'hockey stick' shape of the distribution in Fig.~\\ref{fig:hockey}.  In summary, comparing the observed 2-10keV X-ray and IR luminosities of AGN with estimates of the central black hole mass and host galaxy classification can provide a powerful diagnostic tool to constrain models of AGN activity, $M_{BH}$ and AGN feedback over cosmic time."
    },
    "1005/1005.3716_arXiv.txt": {
        "abstract": "s{ Constraints on the expansion history of the universe from measurements of cosmological distances make predictions for large-scale structure growth. Since these predictions depend on assumptions about dark energy evolution and spatial curvature, they can be used to test general classes of dark energy models by comparing predictions for those models with direct measurements of the growth history. I present predictions from current distance measurements for the growth history of dark energy models including a cosmological constant and quintessence. Although a time-dependent dark energy equation of state significantly weakens predictions for growth from measured distances, for quintessence there is a generic limit on the growth evolution that could be used to falsify the whole class of quintessence models. Understanding the allowed range of growth for dark energy models in the context of general relativity is a crucial step for efforts to distinguish dark energy from modified gravity. } ",
        "introduction": "\\label{sec:intro}  Although several ideas have been proposed to explain the observed acceleration of the cosmic expansion rate, none has yet emerged as a clear favorite from a theoretical viewpoint. Even if we restrict the possibilities to models where general relativity (GR) is valid even on the largest scales and dark energy drives the accelerated expansion, there are still numerous models of dark energy that can fit existing cosmological data. One method for distinguishing among these possibilities is to compare constraints from probes of geometry (distances and the expansion rate) with those from probes of the growth of large-scale structure. Here I present predictions for the growth history from existing measurements of distances and show that while these predictions can vary significantly depending on the specific model of dark energy, there are some generic aspects of the growth predictions that offer the possibility of simultaneously testing large classes of dark energy models.  To maximize the potential for cutting down the allowed space of dark energy models, the goal here is to identify \\emph{general features} of broad classes of dark energy models rather than to place constraints on specific models of dark energy individually. The example of a model class that I will use here is the set of all scalar field quintessence models. A second important point is that for the purposes of this study, the constraints on dark energy parameters themselves are unimportant; instead, the main output of the analysis consists of \\emph{observable predictions} that provide tests of each class of dark energy models. In particular, I will focus on predictions for the growth of large-scale structure.  The growth function describes how initial density fluctuations in the universe grow under the influence of gravity. On large scales, where the density fluctuations $\\delta$ are small enough that the equations for the evolution of perturbations can be linearized, the growth function is independent of scale and can be expressed relative to its value at some redshift $z_{\\rm MD}$ during matter domination as $G(z) = [(1+z)\\delta(z)]/[(1+z_{\\rm MD})\\delta(z_{\\rm MD})]$. During matter domination, $\\delta(z)\\propto (1+z)^{-1}$ so $G(z)=1$, but at late times cosmic acceleration typically causes $G(z)$ to fall below unity. The linear evolution of the growth is related to the Hubble expansion rate $H(z)$ by \\begin{equation} G'' + \\left(4+\\frac{H'}{H}\\right)G' + \\left[ 3+\\frac{H'}{H}-\\frac{3}{2}\\om(z)\\right]G = 0, \\label{eq:growth} \\end{equation} where $\\om(z)=\\om (1+z)^3 [H_0/H(z)]^2$ is the fraction of density in matter with present value $\\om$, and primes denote derivatives with respect to $\\ln a = -\\ln(1+z)$. The Hubble constant, $H_0=H(z=0)$, can also be expressed in the dimensionless form $h = H_0/(100~{\\rm km~s}^{-1}~{\\rm Mpc}^{-1})$.  The expansion rate is constrained observationally through measurements of cosmological distances as a function of redshift, \\begin{equation} D(z) = \\frac{1}{\\sqrt{\\ok} H_0} \\sinh\\left[\\sqrt{\\ok} H_0 \\int_0^z \\frac{dz'}{H(z')} \\right], \\label{eq:distance} \\end{equation} where $\\ok$ parametrizes spatial curvature. Observations of Type Ia supernovae (SNe), the Hubble constant, baryon acoustic oscillations (BAO), and the cosmic microwave background (CMB) all provide constraints on the distance-redshift relation at various redshifts. The inferred evolution of $H(z)$ can be used to predict the linear growth history $G(z)$ using Eq.~(\\ref{eq:growth}). ",
        "conclusions": ""
    },
    "1005/1005.2306_arXiv.txt": {
        "abstract": " ",
        "introduction": "Extra-dimensional theory is an interesting candidate beyond the Standard Model (SM), having rich phenomenology of Kaluza-Klein (KK) particles, a tower of modes for each SM field propagating in bulk. Especially, in the Universal Extra Dimensions (UED) model~\\cite{ued}, the lightest KK particles (LKP) with an odd parity is stable and can be a candidate for the dark matter. The UED can also lower constraints on the compactification scale from the electroweak (EW) precision measurements to few hundred GeV due to the five-dimensional (5D)  Lorentz symmetry \\cite{Appelquist,Gogoladze:2006br}, while the scale is restricted to be larger than few TeV for the brane localized fermion (BLF) scenario \\cite{kk1}-\\cite{Cheung:2001mq}. If the KK-modes are confirmed at the CERN LHC experiment, it would strongly indicate the existence of the extra dimensions.  On the other hand, anther interesting phenomenological consequence of the extra-dimensional theory is the top Yukawa deviation, which is a deviation of the Yukawa coupling between top quark and physical Higgs field from the naive SM expectation: $m_t/v_{\\text{EW}}$. Such a deviation is generally caused when there are several Higgs doublets. Recently it has been pointed out that the deviation can be induced from effects of the brane localized Higgs potentials~\\cite{Haba:2009uu,Haba:2009wa} and the Dirichlet boundary condition (BC)~\\cite{Haba:2009pb} in a simple setup of the extra-dimensional theory with one Higgs doublet.\\footnote{Yukawa deviation is considered in a warped gauge-Higgs unification model~\\cite{Hosotani:2008by}, and a deviation between SM-field and flavon  is considered in~\\cite{Holthausen:2009qj}.}  In this paper, we study a 5D model~\\cite{Haba:2009pb} where the gauge bosons and Higgs doublet of the SM exist in the 5D bulk compactified on a line segment with a flat metric, imposing a generic Dirichlet BC for the Higgs and Neumann BC for all other SM fields. The vacuum expectation value (VEV) of the Higgs is induced from the Dirichlet BC which is generally allowed in higher dimensional theories. We present wave-function profiles of bulk fields under these BCs. The Dirichlet BC removes the Higgs zero-mode and one might worry how a Nambu-Goldstone (NG) boson is supplied for the electroweak symmetry breaking. We show that a superposition of sufficiently large number of KK-modes gives a correct wave-function profile which is flat except at the boundaries.  We consider the case where all other SM fermions propagate in the bulk as the UED model. Then, we will give detailed phenomenology of the model on the top Yukawa deviation, the production and decay of Higgs, constraints on the model from the electroweak precision measurements, and dark matter.  The organization of the paper is as follows. In the next section, we present our extra-dimensional setup and show the profiles, masses and couplings of bulk Higgs and gauge fields. In Section 3, we discuss phenomenological aspects of the setup such as top Yukawa deviation, the production and decay of the Higgs, constraints from the EW precision measurements, and dark matter candidate. In Section 4, we summarize our results. In Appendix we show the 5D gauge and Higgs kinetic Lagrangian and the gauge fixing. ",
        "conclusions": "We have studied the 5D model where the all SM fields exist in the 5D bulk compactified on the line segment with the flat metric. The wave-function profiles of the Higgs and gauge fields have been presented under the Dirichlet BC for the Higgs and the Neumann BC for the gauge fields on the branes. It has been shown that a sufficiently flat profile of the longitudinal component of the zero-mode gauge boson can be obtained by a superposition of the KK-mode of NG bosons.  We have also presented phenomenological discussions on the top Yukawa deviation, the production and decay of Higgs, constraints on the model from the EW precision measurements, and dark matter candidate. In the model, the top Yukawa deviation of 10\\% is predicted. The cross section for the Higgs production, such as the gluon and $WW$ fusions, in this model are over all decreased to 81\\% of the SM expectations. Therefore, the branching ratios are not changed. For the Higgs decay, the width of the process $H\\rightarrow WW$, which is the dominant one, is as large as the Higgs mass. The evaluation of the $S$ and $T$ parameters suggests that the KK scale of $430\\mbox{ GeV}\\lesssim m_H\\lesssim 500\\mbox{ GeV}$ is favored in this scenario.  \\subsection*"
    },
    "1005/1005.0135_arXiv.txt": {
        "abstract": "We model the evolution of buoyant magnetic flux tubes in the Sun's convection zone. A flux tube is assumed to lie initially near the top of the stably stratified radiative core below the convection zone, but a segment of it is perturbed into the convection zone by gradual heating and convective overshoot motions.  The ends (``footpoints'') of the segment remain anchored at the base of the convection zone, and if the segment is sufficiently long, it may be buoyantly unstable, rising through the convection zone in a short time.  The length of the flux tube determines the ratio of buoyancy to magnetic tension: short loops of flux are arrested before reaching the top of the convection zone, while longer loops emerge to erupt through the photosphere.  Using Spruit's convection zone model, we compute the minimum footpoint separation $L_c$ required for erupting flux tubes.  We explore the dependence of $L_c$ on the initial thermal state of the perturbed flux tube segment and on its initial magnetic field strength.  Following an investigation of thermal diffusion time scales and the dynamic rise times of unstable flux tube segments, we conclude that the most likely origin for magnetic flux which erupts to the surface is from short length scale perturbations ($L < L_c$) which are initially stable, but which are subsequently destabilized either by diffusion of heat into the tube or by stretching of the anchor points until $L$ just exceeds $L_c$.  In either case, the separation of the anchor points of the emergent tube should lie between the critical distance for a tube in mechanical equilibrium and one in thermal equilibrium.  Finally, after comparing the dispersion of dynamic rise times with the much shorter observed active region formation time scales, we conclude that active regions form from the emergence of a single flux tube segment. ",
        "introduction": "\\label{section:intro}  In several recent theories of solar and stellar dynamos (\\citealt{Golub1981}; \\citealt{Galloway1981}; \\citealt{DeLuca1986}) magnetic fields are generated not in the convection zone, but in a thin, convectively stable layer just below it.  In order for this magnetic flux to emerge to the photosphere where it can be observed, it must first make its way through the convection zone.  It is generally believed that magnetic buoyancy provides the force which pushes the magnetic flux toward the surface \\citep{Parker1979}. Our goal in this paper is to use the thin flux tube approximation of \\cite{Spruit1981} to study the emergence of buoyant loops of magnetic flux whose ends are still anchored in the stable layer below the convection zone.  Our approach is similar to that used by \\cite{Moreno-Insertis1986}. Further details of our calculations may be found in \\cite{Chou1989}; henceforth Paper I, and \\cite{Fisher1989p}, henceforth Paper II. Here we present only the essentials. ",
        "conclusions": "\\label{section:conclusions}  We have developed a model, based on the thin flux tube approximation of \\cite{Spruit1981}, for studying the emergence of magnetic flux through the convection zone when the footpoints of the flux loop are anchored a distance $L$ apart in the stable layers below.  Figure \\ref{figure:yosfigure1} shows the evolution of a flux tube whose footpoints are sufficiently far apart ($L > L_c$) to allow it to rise to the top of the convection zone.  We have used our model to explore the critical length $L_c$ separating flux tubes which form stable magnetic loops in the convection zone ($L < L_c$) from those which erupt through the photosphere ($L > L_c$).  Figure \\ref{figure:yosfigure2} shows the dependence of $L_c$ on the temperature defect parameter $\\eta$ for two values of the field strength at the base of the convection zone $B_0$, $10^4{\\rm \\ G}$, and $10^6{\\rm \\ G}$.  We find that for perturbations on length scales $L > L_c(B_0,\\eta)$ the rise time is short compared to the duration of the solar cycle.  Under the assumption that the solar cycle can be modeled as a kinematic ($e.g.$, $\\alpha \\omega$ or $\\alpha^2 \\omega$) dynamo model operating below the convection zone, we conclude that such perturbations must therefore be rare or nonexistent.  We speculate that the magnetic flux which does erupt through the photosphere forms initially from perturbations with $L < L_c$, resulting in stable structures.  These are subsequently destabilized either by thermal diffusion or by stretching of the anchor points until $L$ exceeds $L_c$.  In either case, we expect that the anchor point separation $L$ should fall roughly within the range $L_c(\\eta = 0)$ to $L_c(\\eta = 1)$. The rise time of flux tubes is the time scale for conduction of heat into the tube in the first case, and the stretching time in the second case.  Finally, we argue that active regions are formed from the emergence of a single flux tube segment."
    },
    "1005/1005.0245_arXiv.txt": {
        "abstract": "{The stellar helium-to-metal enrichment ratio, $\\Delta Y/ \\Delta Z$, is a widely studied astrophysical quantity. However, its value is still not precisely constrained.} {This paper is focused on the study of the main sources of  uncertainty which affect the $\\Delta Y/ \\Delta Z$ ratio derived from the analysis of the low-main sequence (MS) stars in the solar neighborhood.} {The possibility to infer the value of the helium-to-metal enrichment ratio from the study of low-MS stars relies on the dependence of the stellar luminosity and effective temperature on the initial helium and metal abundances. The $\\Delta Y/\\Delta Z$ ratio is obtained by comparing the magnitude difference between the observed stars and a reference theoretical zero age main sequence (ZAMS) with the related theoretical magnitude differences computed from a new set of stellar models with up-to-date input physics and a fine grid of chemical compositions. A Monte Carlo approach has been used to evaluate the impact on the result of different sources of uncertainty, i.e. observational errors, evolutionary effects, systematic uncertainties of the models. As a check of the procedure, the method has been applied to a different data set, namely the low-MS of the Hyades.} {Once a set of ZAMS and atmosphere models have been chosen, we found that the inferred value of $\\Delta Y/ \\Delta Z$ is sensitive to the age of the stellar sample, even if we restricted the data set to low luminosity stars. The lack of an accurate age estimate of low mass field stars leads to an underestimate of the inferred $\\Delta Y/ \\Delta Z$ of $\\sim 2$ units. On the contrary the method firmly recovers the $\\Delta Y/ \\Delta Z$ value for not evolved samples of stars such as the Hyades low-MS. Adopting a solar calibrated mixing-length parameter and the PHOENIX GAIA v2.6.1 atmospheric models, we found $\\Delta Y/\\Delta Z~= 5.3~\\pm~1.4$ once the age correction has been applied. The Hyades sample provided a perfectly consistent value.  % } {We have demonstrated that the assumption that low-mass stars in the solar neighborhood can be considered as unevolved, does not necessarily hold, and it may indeed lead to a bias in the inferred $\\Delta Y/\\Delta Z$.  The effect of the still poorly constrained efficiency of the superadiabatic convection and of different atmosphere models adopted to transform luminosities and effective temperature into colors and magnitudes have been discussed, too.} ",
        "introduction": "It is a well known and firm result of stellar evolution studies that the main structural, observational and evolutionary characteristics of a star of given mass depend sensitively on the original chemical composition, i.e. the initial helium and metal abundances, $Y$ and $Z$, respectively. As a consequence, these parameters affect also the observable quantities of stellar systems, from star clusters to galaxies.  While the present $Z$ in a stellar atmosphere can be obtained by the direct spectroscopic measurements of some tracer element, mainly iron, with the additional assumption on the mixture of heavy elements, the situation of $Y$ is completely different. In the vast majority of stars the helium lines can not be observed, with the exception of those hotter than $20\\,000$ K. This means that for the low-mass stars, which are the most common and long-lasting objects in the Universe, helium can be directly observed only in advanced evolutionary phases, as the blue part of the horizontal branch or in the post-asymptotic giant branch. Thus, the  actually measured helium abundance is not the original one, but the result of several complex processes, like dredge-up of nuclearly processed material, diffusion and radiative levitation, which severely alter the surface chemical composition.  As a consequence, in order to evaluate the original $Y$, the only possibility is to rely on indirect methods. This explains why such an important parameter is still poorly constrained.  As early suggested by \\citet{peimbert74} from the analysis of the chemical composition of the HII regions in the Large Magellanic Cloud, a common approach in both stellar and population synthesis models is to assume a linear relationship between the original $Y$ and $Z$, \\begin{equation} \\label{eq:dydzlin} Y = Y_\\mathrm{P} + \\frac{\\Delta Y}{\\Delta Z} \\times Z \\quad \\end{equation} where $Y_\\mathrm{P}$ is the primordial helium content, i.e. the result of the Big Bang nucleosynthesis, and $\\Delta Y / \\Delta Z$ the ratio which provides the stellar nucleosynthesis enrichment.  In the last four decades there has been a continuous effort to try to constrain both the $Y_\\mathrm{P}$ \\citep[see e.g.][]{peebles66,churchwell74,peimbert74,peimbert76,lequeux79,kunth83,pagel86,kunth86,pagel92,mathews93, izotov94,olive95,izotov97,olive97,peimbert02,izotov07,peimbert07,spergel07,dunkley09} and the ratio $\\Delta Y / \\Delta Z$ \\citep{faulkner67,perrin77,lequeux79,peimbert80,peimbert86,pagel92,renzini94, fernandes96,pagel98,jimenez03,izotov04,fukugita06,casagrande07}.  Indeed, as previously mentioned, the  relationship between $Y$ and $Z$ adopted in stellar models directly affects some important quantities of stellar systems, both resolved and not, inferred by comparing observations and theoretical predictions. Thus, a precise determination of $Y_\\mathrm{P}$ and $\\Delta Y / \\Delta Z$ is of paramount importance for the studies not only of stellar evolution, but also of galaxy evolution. Furthermore, an accurate estimate of $Y_\\mathrm{P}$ is of great cosmological interest, as it constrains the early evolution of the Universe, when the Big Bang nucleosynthesis occurred. In this paper we will focus on the value of the $\\Delta Y / \\Delta Z$ ratio.  In the past, several techniques have been used to determine such a ratio by means of HII regions, both galactic and extragalactic, planetary nebulae (PNe),  the Sun and chemical evolution models of the Galaxy (see Sect. \\ref{sec:othermeth} for a brief summary of these results).  An alternative and well established way to determine the value of the helium-to-metallicity enrichment ratio takes advantage of the dependence  of the location of stars in the Hertzsprung-Russell (HR) diagram on their helium content. From the study of stellar populations in the galactic bulge \\citet{renzini94} inferred $2 \\leq \\Delta Y / \\Delta Z \\leq 3$. A frequently adopted approach relies on the analysis of the fine structure of the low-MS of the local field stars in the HR diagram. Pioneers of such an approach have been Faulkner, who found  $\\Delta Y / \\Delta Z = 3.5$  \\citep{faulkner67} and Perrin and collaborators, who obtained $\\Delta Y / \\Delta Z = 5$ \\citep{perrin77}. Following these early studies, \\citet{fernandes96} constrained the value of $\\Delta Y / \\Delta Z$ to be larger than 2, by comparing the broadening in the HR diagram of the low-mass MS stars in the solar neighborhood and the theoretical ZAMS of several $Y$ and $Z$. With a similar approach but taking advantage of Hipparcos data, \\citet{pagel98} inferred $\\Delta Y / \\Delta Z = 3 \\pm 2 $. This kind of approach culminated recently in the works by \\citet{jimenez03} and  \\citet{casagrande07} who provided, respectively, $\\Delta Y / \\Delta Z = 2.1 \\pm 0.4 $  and 2.1 $\\pm$ 0.9.  The present analysis deals with the determination of the $\\Delta Y / \\Delta Z$ ratio by means of the comparison between the local K dwarf stars, for which accurate measures of both the [Fe/H] and parallaxes are available, and state-of-the-art stellar models. A great effort has been devoted to discuss the effect of the main uncertainties still present in stellar models on the inferred value of $\\Delta Y / \\Delta Z$.  In Sect. \\ref{sec:models} we present the set of low-mass stellar models we have computed for this paper; in Sect. \\ref{sec:dataset} we describe the data set we have used; Section \\ref{sec:method} contains the description of the analysis method, while Sect. \\ref{sec:uncert} deals with the possible sources of uncertainty that could affect the method itself. Results for the adopted data set are presented in Sect. \\ref{sec:results}. In Sect. \\ref{sec:nonlin} we investigate the possibility of a non linear relation between $Y$ and $Z$. In Sect. \\ref{sec:Hyatest} we apply our method to an independent and unevolved set of stars, i.e. the Hyades low-main sequence. We compare our results to those obtained by other authors, with independent methods, in Sect. \\ref{sec:othermeth}. Section \\ref{sec:concl} contains the final discussion and summary of the whole paper. ",
        "conclusions": "\\label{sec:concl}  The principal aim of this work was to test the reliability of the determination of $\\Delta Y/ \\Delta Z$ by the comparison between low-MS stars and theoretical ZAMS models. A very fine grid of stellar models has been computed for many values of $\\Delta Y/ \\Delta Z$, [Fe/H] and masses adopting two different mixing-length parameters $\\alpha$, namely 1.97 (our solar calibration) and 2.4, and two different sets of atmosphere models to transform luminosities and effective temperatures into magnitudes and color indices, the PHOENIX GAIA v2.6.1 models \\citep{brott} and the ATLAS9 ones \\citep{castelli03}, respectively.  A detailed analysis of the capabilities of the method and of the main uncertainty sources affecting the derived results has been performed by means of many numerical experiments on synthetic data set produced under controlled conditions and with precisely known properties.  One of the main findings of the paper is that the inferred value of $\\Delta Y/ \\Delta Z$ is quite sensitive to the age of the stellar sample, even in the case in which only very faint (i.e. $M_V > 6$ mag) MS stars are selected. By means of numerical experiment we showed that the lack of an age estimate of low mass field stars leads to an underestimate of the inferred $\\Delta Y/ \\Delta Z$ of about $2$ units. As a consequence, the face value of the helium-to-metals enrichment ratio provided by the recovery procedure applied to the solar neighborhood stellar sample must be corrected for this age-bias.  Adopting our reference set of models (i.e. those with our ``solar'' calibrated $\\alpha$) transformed into the observational plane using the PHOENIX GAIA v2.6.1 model atmospheres, we found $\\Delta Y/\\Delta Z = 5.3 \\pm 1.4$.  Such a result has been checked against an independent and very accurate data set, that is, the low-MS of the Hyades cluster. This sample has the additional advantage to be unaffected by the age-bias, since its very young age (i.e. $500 \\div 600$ Myr) guarantees that the low-mass ($<0.9 M_{\\sun}$) MS stars which constitutes the sample are essentially unevolved. The recovery method provided again $\\Delta Y/\\Delta Z = 4.75 \\pm 0.35$, in perfect agreement with the result for field stars.  To further check the consistency of this final result, we calculated isochrones with $\\Delta Y/ \\Delta Z = 5$ (the closest in our grid of models) and the measured [Fe/H] of Hyades. The good agreement between these isochrones and the Hyades MS, not only in the faint part belonging to the sample but also at higher luminosities, is a further proof of the internal consistency of the recovery procedure.  The effect of a change in the assumed efficiency of the superadiabatic convection (i.e. the $\\alpha$ parameter) in the stellar models used to build the ZAMS and in the adopted atmosphere models used to transform luminosities and effectives temperatures into magnitudes and color indices has been discussed, too. More in detail, the recovery method yield a nominal value of $\\Delta Y/\\Delta Z=  1.59 \\pm 1.01$ when adopting the set of theoretical models computed with $\\alpha=~2.4$ and transformed into the observational plane by means of PHOENIX atmosphere models and $\\Delta Y/\\Delta Z= 0.5$ when adopting our standard set of models with $\\alpha=~1.97$ but the ATLAS9 atmosphere models. These values become about 3.6 and 2.5, respectively, once corrected for the age bias.  Our data have [Fe/H] determinations coming from two different sources, i.e. \\cite{Nordstrom04} and \\cite{Taylor05} catalogs. The 4 stars we have in common from the two catalogs show a disagreement in the [Fe/H] scale between the two catalogs. This disagreement was already identified in \\cite{Taylor05} and the author provides a Table to account for it and put the two catalogs on the same [Fe/H] scale. Nevertheless, even after the re-zeroing procedure suggested by \\cite{Taylor05}, we still find some disagreement in the final results on $\\Delta Y/ \\Delta Z$ when the two catalogs are considered separately. Probably an even more accurate study of the zero points of metallicity determinations is needed."
    },
    "1005/1005.2847_arXiv.txt": {
        "abstract": "Terzan~5 is a globular cluster-like stellar system in the Galactic Bulge which has been recently found to harbor two stellar populations with different iron content and probably different ages \\citep{fe09}. This discovery suggests that Terzan~5 may be the relic of a primordial building block which contributed to the formation of the Galactic Bulge. Here we present a re-determination of the structural parameters (center of gravity, density and surface brightness profiles, total luminosity and mass) of Terzan~5, as obtained from the combination of high-resolution ({\\it ESO}-MAD and {\\it HST} ACS-WFC) and wide-field ({\\it ESO}-WFI) observations. We find that Terzan~5 is significantly less concentrated and more massive than previously thought. Still it has the largest collision rate of any stellar aggregate in the Galaxy.  We discuss the impact of these findings on the exceptional population of millisecond pulsars harbored in this stellar system. ",
        "introduction": "Terzan~5 is commonly catalogued as a globular cluster (GC) located in the inner Bulge of our Galaxy.  It is difficult to observe because it is heavily reddened, with an average color excess $E(B-V)=2.38$ (Barbuy et al. 1998; Valenti et al. 2007). Not only is the reddening large, but it strongly depends on the line of sight (differential reddening; see Ortolani et al. 1996, Valenti et al. 2007).  Terzan~5 has an exceptionally large population of millisecond pulsars (MSPs). Indeed, the 33 MSPs detected so-far in Terzan~5 amount to about $25\\%$ of the entire sample of known MSPs in Galactic GCs (Ransom et al. 2005; see the updated list at {\\tt http://www.naic.edu/pfreire/GCpsr.html})  As part of a project (Ferraro et al. 2001, 2003; Cocozza et al. 2008) aimed at studying the properties of stellar populations harboring MSPs, we obtained a set of high-resolution images of Terzan~5 in the $K$ and $J$ bands using the multi-conjugate adaptive optics (AO) demonstrator MAD \\citep{marchetti07} temporally installed at the European Southern Observatory (ESO) Very Large Telescope (VLT).  The ($K,~J-K$) color-magnitude diagram (CMD) obtained from these observations led to the discovery of two well-defined red horizontal branch (HB) clumps, clearly separated in luminosity ($\\delta K\\sim 0.3$) and color (see Ferraro et al. 2009, hereafter F09; also see Figure \\ref{fig:cmd}). A prompt spectroscopic follow-up demonstrated that the two populations have the same radial velocity (hence they belong to the same stellar systems) and their metal content is different: ${\\rm [Fe/H]} \\simeq -0.2$ and ${\\rm [Fe/H]}\\simeq +0.3$ for the fainter and the brighter group, respectively. These findings and the comparison with theoretical stellar isochrones confirm the existence of two distinct stellar populations in Terzan~5 and suggest that they possibly have been generated by two bursts of star formation with a time separation of a few ($\\approx 6$) Gyrs.  While the age gap can be reduced by invoking a difference in the helium content of the two populations \\citep{dantona10}, the iron enrichment and the spatial segregation of the brightest clump, together with the extraordinary amount of MSPs found in Terzan 5, indicate that this system probably experienced a particularly troubled formation and evolutionary history (see Sect. 5).  Terzan~5 is the first a GC-like system in the Galactic Bulge found to have a spread in the iron content and it could be the relic of one of the building blocks that contributed to the formation of the Bulge.  Indeed the discovery might represent the observational evidence that even the innermost part of galactic spheroids form (at least partially) by the accretion/merging of small, previously formed and internally evolved stellar systems \\citep[e.g.,][]{immeli04}.  In this paper we present the accurate re-determination of Terzan~5 structural parameters (surface density and surface brightness profiles, total luminosity, collision rate, etc.), obtained from a combination of high-resolution and wide-field observational data. These parameters provide basic information for a deeper understanding of the origin and the evolution of this puzzling system. ",
        "conclusions": "The star density and SB profiles can be used to derive the integrated luminosity of the cluster. From the best-fit King model we estimate that the percentage of cluster light within regions of radius $r=15\\arcsec$, $18\\arcsec$, and $20\\arcsec$ are roughly 30\\%, 36\\% and 40\\%, respectively.  Using aperture photometry on the MAD images, we obtain integrated-light values of $K(r<15\\arcsec) = 3.44$, $K(r<18\\arcsec) = 3.3$ and $K(r<20\\arcsec) = 3.2$ mag, respectively. Adopting the color excess $E(B-V)=2.38$, the distance modulus $(m-M)_0=13.87$ (Valenti et al. 2007, corresponding to a distance of $d=5.9\\pm0.5$ kpc) and the bolometric correction $BC_K = 2.4$ appropriate for a population of intrinsic color $(J-K)_0 = 0.8$ (see Montegriffo et al. 1998), we estimate that the corresponding bolometric luminosity in the considered regions is: $L_{\\rm bol} (r < 15\\arcsec) = 3 \\times 10^5 L_\\odot$, $L_{\\rm bol} (r < 18\\arcsec) = 3.4 \\times 10^5 L_\\odot$ and $L_{\\rm bol} (r < 20\\arcsec) = 3.7 \\times 10^5 L_\\odot$.  Considering the fraction of light sampled in each region we find that the total luminosity of the system is $L_{\\rm bol} = 9.5\\pm0.3 \\times 10^5 L_\\odot$.  An independent estimate of the total luminosity of the stellar system can be derived from its stellar population, by using a simple relation \\citep{fuel86} linking the number of stars ($N_j$) observed in a given post-main sequence evolutionary stage $j$ and the luminosity of the entire parent cluster ($L_T$): \\begin{equation} N_j = B \\times t_j \\times L_T, \\end{equation} where $B$ is the specific evolutionary flux (for intermediate/old stellar populations $B = 2 \\times 10^{-11}\\,{\\rm stars\\, yr}^{-1} L_\\odot^{-1}$) and $t_j$ is the duration of the evolutionary stage. The number of HB stars counted in the two clumps by F09 is quite large: a total of about 1300 (with 800 and 500 belonging to the faint and the bright HB clumps, respectively). This population is comparable to (or even larger than) that observed in the largest Galactic GCs, like 47 Tucanae \\citep{bec_47tuc} and NGC~6388 \\citep{ema_6388}, and suggests that the overall size of Terzan 5 (in terms of luminosity and mass) is comparable to that of these systems. For a quantitative estimate, we insert the observed number of HB stars in the above relation and adopt $t_{\\rm HB} = 10^8$ yr. This provides a luminosity of $4 \\times 10^5 L_\\odot$ and $2.5 \\times 10^5 L_\\odot$ for the two parent populations, and a total luminosity of $6.5 \\times 10^5 L_{\\odot}$ for the entire stellar system.  This estimate, which is distance and reddening independent, is quite consistent with the previous one, thus confirming that Terzan~5 has a considerable total luminosity (hereafter we adopt the average value $L_{\\rm bol} = 8 \\times 10^5 L_\\odot$), significantly higher than previously thought. By comparison, adopting the values of distance and reddening quoted above and a bolometric correction $L_{\\rm bol}\\simeq 1.4\\,L_V$, the total bolometric luminosity corresponding to the integrated magnitude ($V_t=13.85$) quoted by \\citet{harris96} would be only $L_{\\rm bol}\\simeq 10^5 L_\\odot$.  The discrepancy is most probably due to the strong (differential) reddening affecting the system, especially in the optical bands. This effect is greatly reduced for our new estimate, since it is based on the observed $K$-band integrated magnitude and the number of HB stars.  By assuming a mass-to-light ratio $M/L_{\\rm bol}=3$ \\citep[e.g.,][]{maraston98}, the total stellar mass of this system is $M_{\\rm T} \\simeq 2 \\times 10^6 M_\\odot$.  \\citet{verhut87} first suggested that the collision rate of Terzan~5 is the highest among the Galactic GCs. We can now re-compute this quantity by adopting the newly determined parameters. Following \\citet{verhut87}, the collisional parameter ($\\Gamma$) for a King-model or virialized system can be computed as: $\\Gamma \\propto \\rho_0 \\times r_c^{0.5}$, where $\\rho_0$ is the central mass density. By using the values obtained above and equation (7) of \\citet{djorg93}, we find that the collision parameter of Terzan~5 is between 5 and 10 times higher than that of Liller 1 and of other massive clusters for which structural parameters have been recently re-determined \\citep[NGC6388, NGC6266, 47Tuc;][respectively]{ema_6388, bec_6266, mapelli06}. Hence we confirm that, even with the new structural parameters (suggesting a lower concentration and a larger mass than previously thought), Terzan~5 still has the largest known collision rate of any stellar aggregate in the Galaxy.  The co-existence of two stellar populations with different iron content (and probably ages) suggests that the original mass of Terzan~5 was significantly larger in the past than observed today, large enough to retain the iron-enriched gas that, otherwise, would have been ejected out from the system by the violent supernova (SN) explosions.  Indeed, the smallest systems with solid evidences of a spread in the iron content (and ages) are significantly more massive than GCs: the dwarf spheroidal satellites of the Milky Way typically have masses of $\\sim 10^7 M_\\odot$ \\citep[Strigari et al. 2008; see also][]{battaglia08} and, following recent chemo-dynamical models well reproducing the observations, their initial masses amounted to a few $10^8 M_\\odot$ \\citep{revaz09}. While a lower limit of $\\sim 10^7 M_\\odot$ for the proto-Terzan 5 could also be hazarded following \\citet{baumgardt08}, more detailed and extensive simulations are needed to firmly determine the smallest total mass necessary to retain the SN ejecta.  The exceptionally high metallicity regime of the two stellar populations found in Terzan~5 also suggests a quite efficient enrichment process, that could have a relevant role in the origin of its population of MSPs. In particular, both the iron and the [$\\alpha/$Fe] abundance ratios measured in Terzan~5 \\citep[][Rich et al. 2010, in preparation]{origlia04} show a remarkable similarity with those of the Bulge stars.  This strongly suggests that these two structures shared the same star formation and chemical enrichment processes.  The many observations of Bulge stars \\citep[e.g.,][and references therein]{melendez08,origlia08,ryde09} indicate that they are all characterized by an old age, a high (close to solar) average metallicity [Fe/H], and an [$\\alpha/$Fe] ratio which is enhanced (due to SNII enrichment) up to a metallicity [Fe/H]$\\simeq 0$. These constraints suggest a scenario where the dominant stellar population of the Bulge formed early (thus explaining the old age\\footnote{\\scriptsize{Additional episodes of star formation mainly confined in the innermost ($\\sim 100$ pc) region could eventually explain the presence of younger stars \\citep[e.g.,][]{blum03,figer04}}.}), rapidly and with high efficiency (from a gas mainly enriched by SNII, thus explaining the [$\\alpha/$Fe] enhancement up to high iron contents\\footnote{\\scriptsize{The [$\\alpha/$Fe]--[Fe/H] relation shows a down-turn at a value of [Fe/H] which depends on the star formation rate: the higher the latter, the higher the metallicity at which the down-turn occurs. Such a value is [Fe/H]$\\simeq -1$ in the Old Halo/Disk, while it is significantly higher ([Fe/H]$\\simeq 0$) in the Bulge, testifying a much higher star formation rate in this dense environment.}}).  Also chemical evolution models \\citep[e.g.,][]{ballero07,mcwilliam08} indicate that the abundance patterns observed in the Bulge require a quite high star formation efficiency and an initial mass function flatter than that in the solar neighbourhood, thus to rapidly enrich the gas up to about solar metallicity through an exceptionally large amount of SNII explosions. The assumption of a similar scenario for Terzan~5 would naturally explain its extraordinary population of MSPs, since the expected high number of SNII would produce a large population of neutron stars, most of which would have been retained by the deep potential well of the massive proto-Terzan~5 system.  Then the high collision rate could have favoured the formation of binary systems containing neutron stars and promoted the re-cycling process that finally generated the large population of MSPs now observed in Terzan~5. If such a scenario is correct, many more MSPs still wait to be discovered in this system \\citep[see also][]{ransom05}, the 33 known objects probably being just the tip of the iceberg. Future deeper pulsar searches of Terzan~5, perhaps with larger telescopes such as the Square Kilometer Array, will shed additional light on the nature of this system."
    },
    "1005/1005.0842_arXiv.txt": {
        "abstract": "We present a comprehensive spectral analysis of all \\integral\\/ data obtained so far for the X-ray--bright Seyfert galaxy NGC 4151. We also use all contemporaneous data from \\xte, \\xmm, \\swift\\/ and \\suzaku. We find a linear correlation between the medium and hard-energy X-ray fluxes measured by \\integral, which indicates an almost constant spectral index over six years. The majority of \\integral\\/ observations were made when the source was either at a very bright or very dim hard--X-ray state. We find that thermal Comptonization models applied to the bright state yields the plasma temperature of $\\simeq 50$--70 keV and its optical depth of $\\simeq 1.3$--2.6, depending on the assumed source geometry. For the dim state, these parameters are in the ranges of $\\simeq 180$--230 keV and $\\simeq 0.3$--0.7, respectively. The Compton parameter is $y\\simeq 1$ for all the spectra, indicating a stable geometry. Using this result, we can determine the reflection effective solid angles associated with the close and distant reprocessing media as $\\simeq 0.3\\times 2\\upi$ and $0.2 \\times 2\\upi$, respectively. The plasma energy balance, the weak disc reflection and a comparison of the UV fluxes illuminating the plasma to the observed ones are all consistent with an inner hot accretion surrounded by an outer cold disc. The disc truncation radius can be determined from an approximate equipartition between the observed UV and X-ray emission, and from the fitted disc blackbody model, as $\\sim 15$ gravitational radii. Alternatively, our results can be explained by a mildly relativistic coronal outflow. ",
        "introduction": "\\label{intro}  Continuum properties of the hard X-ray and soft $\\gamma$-ray emission from radio-quiet Seyfert galaxies nuclei are relatively well known thanks to the many satellites operating during the last decades. The spectra are commonly approximated by a phenomenological e-folded power-law model, $F(E)\\propto E^{1-\\Gamma} \\exp(-E/E_{\\rm c})$ (where $\\Gamma$ is the photon index and $E_{\\rm c}$ is the e-folding, or cut-off, energy), accompanied by a Compton reflection component. A recent study based on a local ($z < 0.1$) sample of 105 Seyfert galaxies observed with \\sax\\/ presents the average parameters of this model for Seyfert 1s and 2s \\citep{Dadina2008}. The mean $\\Gamma$ is $\\approx 1.9$ (1.8) for Seyfert 1 (2) nuclei, $E_{\\rm c}$ is $\\simeq 230$ (380) keV, and the average relative strength of reflection, $R\\equiv \\Omega/2\\upi$ (where $\\Omega$ is the effective solid angle subtended by the reflector), is $\\simeq 1.2$ (0.9). See \\citet{Beckmann2009} for the similar \\integral\\/ results.  This phenomenological model is too crude for studying of the physics of the central engine of Seyferts. There is a general consensus that their X$\\gamma$ photons are from (predominantly thermal) Comptonization by hot electrons of some seed soft photons. The seed photons may come from an optically-thick accretion disc or clouds in the vicinity of the hot plasma, or may be internally produced by the synchrotron process, e.g., \\citet{Xie2010}. For example, non-simultaneous data for the Seyfert NGC 5548 from \\rosat, \\ginga\\/ and \\gro/OSSE have yielded the plasma temperature, $kT_{\\rm e}\\simeq 55$ keV, and the Thomson optical depth, $\\tau\\simeq 2$ \\citep{Magdziarz1998}. The average OSSE spectra of 11 Sy 1s and 8 Sy 2s have yielded similar parameters, $kT_{\\rm e}\\simeq 70$--80 keV, $\\tau\\simeq 1.7$ \\citep*{Zdziarski2000}. On the other hand, \\citet{Petrucci2000} obtained a much higher $kT_{\\rm e}\\simeq 250$ keV (and $\\tau\\simeq 0.2$--0.4) for NGC 5548 using \\sax\\/ data and a different anisotropic Comptonization model in the slab geometry. Then, \\citet{Petrucci2001a}, hereafter P01, found, with the same model, a similarly high values of $kT_{\\rm e}\\simeq 170$--320 keV (and $\\tau\\simeq 0.05$--0.20) for 6 Sy 1s observed by \\sax. We note, however, that such high $kT_{\\rm e}$ yield spectra being well above the OSSE fluxes at $\\ga 100$ keV for either NGC 5548, NGC 4151 or the average Seyferts.  The nearby, $z = 0.0033$, Seyfert 1.5 galaxy NGC 4151 is the second (after Cen A) brightest persistent AGN in the 20--100 keV band. Its electromagnetic spectrum has been extensively studied, with its properties in the radio, infrared, optical, ultraviolet and X$\\gamma$ bands well established. Its black-hole mass has been estimated based on reverberation as $(4.6^{+0.6}_{-0.5}) \\times 10^{7}\\msun$ \\citep{Bentz2006}. Although this estimate is subject to a systematic uncertainty by a factor of $\\sim$3--4, it is consistent with the estimate based on scaling of the power spectrum to Cyg X-1 of \\citet{Czerny2001}. The overall picture of the region surrounding the nucleus, based on {\\it HST}/STIS, \\chandrae and \\xmme data, is that the complex absorption observed for NGC 4151 is due to massive outflows, presumably disc winds, forming several regions characterized by a range of column densities, $N_{\\rm H}$, and ionization levels \\citep{Schurch2004,Kraemer2006}. This explains earlier results from \\rosat, \\asca, \\ginga\\/ and \\xte\\/ satellites, with the medium-energy X-ray spectra fitted only after applying a complex absorber model, consisting of several components fully and partially covering the central source (e.g., \\citealt*{Zdziarski1996a,Zdziarski2002}, hereafter Z02). The absorber undergoes rapid daily changes of $N_{\\rm H}$ \\citep{Puccetti2007}. Also, a narrow Fe $K_{\\alpha}$ line with the equivalent width of 50--200 eV is observed (Z02; \\citealt{Schurch2003,DeRosa2007}).  NGC 4151 has frequently been observed by \\gro\\/ and \\sax. The e-folded power law model yields $E_{\\rm c}\\simeq 50$--200 keV \\citep{Johnson1997,DeRosa2007}. Thermal Compton models yield $kT_{\\rm e}\\simeq 50$--80 keV, $\\tau\\simeq 1$--2 (\\citealt{Johnson1997}; Z02). The strength of Compton reflection is moderate, $R \\approx 0.4$ (e.g., \\citealt{DeRosa2007}).  NGC 4151 was observed by \\integral\\/ 8 times as a primary or secondary target. The results of the first dedicated observation of 2003 May are given in \\citet{Beckmann2005}. The object was found in a spectral state similar to those observed with OSSE, but at the highest flux level ever noticed. There was no significant spectral variability despite $\\sim 50$ per cent variations of the flux. A thermal Compton model fitted to the summed spectrum yielded $kT_{\\rm e}\\simeq 90$ keV, $\\tau\\simeq 1.3$, $R\\simeq 0.7$, $N_{\\rm H}\\simeq 7 \\times 10^{22}$ cm$^{-2}$. During almost all later \\integral\\/ observations the source was found at moderate to very low flux levels. This has allowed us to study the spectral properties of the dim state of this AGN with unprecedented precision and compare them with the bright state of 2003 May.  \\section[]{Observations and Data Reduction} \\label{data}  We use all NGC 4151 data collected by \\integral\\/ as of 2010 January, see Table \\ref{obsint}. Good quality data come from the dedicated NGC 4151 observations made in 2003 May, 2007 January, May and December, and 2008 May. The remaining data are taken from observations of the Coma cluster, Mrk 273, NGC 4736, M51 and Mrk 421, when NGC 4151 was almost always seen at an off-axis angle of 9\\degr--15\\degr. We select the data with the off-axis angle $<15\\degr$ for ISGRI and SPI, and $<3\\degr$ for the JEM-X and OMC. The data have been reduced using the Offline Scientific Analysis ({\\sc osa}) 7.0 provided by the \\integral\\/ Science Data Centre \\citep{Courvoisier2003}, with the pipeline parameters set to the default values. Since the {\\sc osa} v.\\ 8.0 of 2009 August brought a major improvement for the JEM-X data analysis, we have used it for these data. The ISGRI and SPI spectra and light curves have been extracted with the standard spectral extraction software, including the catalogue sources NGC 4151, NGC 4051, NGC 4138, Mrk 766, NGC 4258, NGC 4395, Coma Cluster, NGC 5033, Mrk 421 and Mrk 268. The JEM-X spectra have been obtained from mosaic images stacked for a given observation. We use the standard response files for all instruments. The OMC magnitude was converted into flux using the calibration of \\citet{Johnson1966}, $F({\\rm V}\\!=\\!0)= 3.92\\times 10^{-9}$ erg cm$^{-2}$ s$^{-1}$ \\AA$^{-1}$.  \\begin{table*} \\centering \\caption{The observation log of \\integral. The ISGRI exposure is effective, corresponding to fully coded observations, and the JEM-X one is summed over all pointings with the off-axis angle $<3\\degr$. The last column gives the spectral set as defined in Section \\ref{assumptions}; B = bright, D = dim, M = medium.} \\label{obsint} \\begin{tabular}{@{}cccccccc@{}} \\hline Revolution & Start time, UTC (MJD)  & End time, UTC (MJD)  & Eff. exposure [s] & JEM-X exposure [s] & Spectral set \\\\ \\hline 0036 & 2003-01-29 16:43 (52668.697) & 2003-01-31 08:18 (52670.346) &  11296 & -- & -- \\\\ 0071 & 2003-05-14 12:42 (52773.529) & 2003-05-16 11:39 (52775.485) &   9908 & -- & -- \\\\ 0072 & 2003-05-17 13:22 (52776.557) & 2003-05-18 19:56 (52777.831) &   5376 & -- & -- \\\\ 0073 & 2003-05-20 18:44 (52779.781) & 2003-05-22 10:23 (52781.433) &  20441 & -- & -- \\\\ 0074 & 2003-05-23 08:35 (52782.358) & 2003-05-25 11:42 (52784.488) & 117528 &  91633 & B \\\\ 0075 & 2003-05-25 21:03 (52784.877) & 2003-05-28 05:33 (52787.231) & 153481 & 153481 & B \\\\ 0076 & 2003-05-28 20:51 (52787.869) & 2003-05-29 12:13 (52788.509) &  39533 &  39533 & -- \\\\ 0274 & 2005-01-10 16:05 (53380.670) & 2005-01-12 17:36 (53382.733) &   9146 & -- & -- \\\\ 0275 & 2005-01-13 19:36 (53383.817) & 2005-01-15 04:20 (53385.181) &  11745 & -- & -- \\\\ 0310 & 2005-04-29 16:32 (53489.689) & 2005-04-30 03:46 (53490.157) &   1734 & -- & D \\\\ 0311 & 2005-05-01 08:54 (53491.371) & 2005-05-02 20:28 (53492.853) &   3290 & -- & D \\\\ 0312 & 2005-05-03 22:44 (53493.947) & 2005-05-04 09:57 (53494.416) &   1062 & -- & D \\\\ 0317 & 2005-05-19 18:04 (53509.753) & 2005-05-21 01:09 (53511.048) &  10144 & -- & D \\\\ 0318 & 2005-05-22 01:03 (53512.044) & 2005-05-24 06:37 (53514.276) &   7894 & -- & D \\\\ 0324 & 2005-06-09 17:55 (53530.747) & 2005-06-10 23:12 (53531.967) &   9748 & -- & D \\\\ 0448 & 2006-06-14 09:54 (53900.413) & 2006-06-17 00:03 (53903.002) &  29525 & -- & D \\\\ 0449 & 2006-06-17 10:19 (53903.430) & 2006-06-19 23:47 (53905.991) &  28497 & -- & D \\\\ 0450 & 2006-06-20 09:23 (53906.391) & 2006-06-22 23:29 (53908.978) &  26974 & -- & D \\\\ 0451 & 2006-06-23 09:06 (53909.379) & 2006-06-25 23:10 (53911.965) &  26844 & -- & D \\\\ 0521 & 2007-01-18 16:45 (54118.698) & 2007-01-20 23:28 (54120.978) & 122878 & 122878 & D \\\\ 0522 & 2007-01-21 16:35 (54121.691) & 2007-01-24 06:07 (54124.255) & 123751 & 105848 & D \\\\ 0561 & 2007-05-18 09:31 (54238.397) & 2007-05-20 16:11 (54240.674) & 110696 & 110696 & D \\\\ 0562 & 2007-05-22 00:44 (54242.031) & 2007-05-23 21:34 (54243.899) & 100729 &  78177 & D \\\\ 0563 & 2007-05-24 09:06 (54244.379) & 2007-05-25 06:31 (54245.272) &  46903 &  46903 & D \\\\ 0634 & 2007-12-22 17:53 (54456.745) & 2007-12-25 07:19 (54459.305) & 131811 & 131811 & M \\\\ 0636 & 2007-12-28 17:26 (54462.726) & 2007-12-31 01:53 (54465.078) & 119738 & 100028 & M \\\\ 0678 & 2008-05-02 21:17 (54588.887) & 2008-05-04 22:36 (54590.942) &  24805 & -- & M \\\\ 0679 & 2008-05-05 20:47 (54591.866) & 2008-05-07 07:09 (54593.298) &  17266 & -- & M \\\\ 0809 & 2009-05-31 11:31 (54982.480) & 2009-05-31 20:58 (54982.874) &  20019 & 2891 & -- \\\\ 0810 & 2009-06-01 05:49 (54983.242) & 2009-06-03 16:10 (54985.674) & 112086 & 46546 & -- \\\\ 0811 & 2009-06-04 05:37 (54986.234) & 2009-06-06 10:24 (54988.433) & 108501 & 40239 & -- \\\\ \\hline \\end{tabular} \\end{table*}  To better constrain the spectra at low energies, we supplemented the \\integral\\/ data by all available X-ray observations of NGC 4151 taken since 2003 by \\xte, \\xmm\\/ and \\suzaku, see Table \\ref{obsother}. The \\xte\\/\\/ PCU2 light curves and spectra and HEXTE spectra were extracted with the {\\sc heasoft} 6.5.1, using standard selection criteria. The 5 \\xmm\\/ observations (denoted here X1--X5) were analyzed with {\\sc sas} 8.0.1. We used only the EPIC pn data, and excluded periods of high or unstable background. The spectra X1 and X2 have been found to be compatible with each other; we have therefore added them together. For the single \\suzaku\\/ observation, the data in the $3\\times 3$ mode were reduced with {\\sc heasoft} 6.6. The spectra of the front-illuminated CCD, XIS0 and XIS3, were added together, whereas the spectrum of the back-illuminated detector XIS1 was used separately. We also extracted the HXD/PIN spectrum using the standard procedure. Although there have been a number of \\chandra\\/ NGC 4151 observations, all those without gratings that were public at the time we started the analysis were from before 2000 March.  Tables \\ref{obsint}--\\ref{obsother} identify the data used for our three main spectral sets, bright (B), medium (M), and dim (D), whose selection is based on ISGRI flux. The remaining \\integral\\/ data, not marked with a letter in Table \\ref{obsint}, were not used for spectral analysis. We have excluded them because either there were no corresponding high-quality spectra from X-ray satellites (e.g., Revs.\\ 0809--0811) and/or their flux level was outside the range assumed by us (e.g., beginning of Rev.\\ 0074 and Rev.\\ 0076 in the case of the B state). The data sets from the other satellites were not always simultaneous with the \\integral\\/ data. Thus, we assign them to one of the three sets based again on the flux, using the monitoring data from \\swift/BAT (see Sections \\ref{variability}, \\ref{assumptions} for details of the selection).  \\begin{table*} \\centering \\caption{The observation log for the other X-ray satellites. The exposure times are given for the \\xte\\/ PCU2 (where the number of pointings added together is shown in parentheses), EPIC pn (\\xmm) and XIS (\\suzaku) detectors. The last column shows the data set as defined in Table \\ref{obsint}, and, in the case of \\xmm, its consecutive observation number.} \\label{obsother} \\begin{tabular}{@{}lcccccr@{}} \\hline Obs. ID & Start time, UTC (MJD) & End time, UTC (MJD) & Exposure & Spectral set \\\\ \\hline \\multicolumn{5}{c}{\\it RXTE} \\\\ 80416-01-01-(00--09) & 2003-05-24 05:49 (52783.242) & 2003-05-29 04:18 (52788.179) & 15760 (11) & B \\\\ 92113-08-(06--09)    & 2006-05-13 02:28 (53868.103) & 2006-06-23 15:42 (53909.654) & 4240 (4)   & D \\\\ 92113-08-(22--35)    & 2006-12-23 12:51 (54092.535) & 2007-06-22 08:57 (54273.373) & 14272 (14) & D \\\\ \\hline \\multicolumn{5}{c}{\\it XMM-Newton} \\\\ 0143500101 & 2003-05-25 01:38 (52784.068) & 2003-05-25 06:54 (52784.288) &  5821 & X1, B \\\\ 0143500201 & 2003-05-26 20:35 (52785.858) & 2003-05-27 01:51 (52786.077) & 11389 & X2, B \\\\ 0143500301 & 2003-05-27 15:17 (52786.637) & 2003-05-27 20:33 (52786.856) & 12291 & X3, B \\\\ 0402660101 & 2006-05-16 06:22 (53871.265) & 2006-05-16 17:35 (53871.733) & 27980 & X4, D \\\\ 0402660201 & 2006-11-29 17:20 (54068.722) & 2006-11-30 08:00 (54069.333) & 21468 & X5, M \\\\ \\hline \\multicolumn{5}{c}{\\it Suzaku} \\\\ 701034010 & 2006-12-18 20:05 (54087.837) & 2006-12-21 09:14 (54090.385) & 124980 & D \\\\ \\hline \\end{tabular} \\end{table*}  \\section[]{Variability} \\label{variability}  NGC 4151 has been observed with X$\\gamma$ detectors for almost 40 years. The satellite and balloon results up to 1988 were compiled by \\citet{Perotti1991}. Then, it has been observed above 20 keV by \\ginga, \\granat, \\gro, \\sax, \\integral, \\swifte and \\suzaku. Fig.\\ \\ref{longest}(a) shows the 20--100 keV flux observed since 1972 October compared to the optical flux and the \\xte/ASM count rate. Because the spectra from earlier observations are not publicly available, the 20--100 keV fluxes were determined using the values of the flux at 35 keV and of the photon index presented in table 2 of \\citet{Perotti1991}. For this reason, we do not show their uncertainties. Similarly, the \\granat\\/ fluxes were determined using table 3a of \\citet{Finoguenov1995}, where the ART-P and SIGMA spectra are fitted by a power-law model. For all other observations, we use the spectra from {\\sc heasarc} (\\gro/OSSE, \\sax/PDS) or spectra extracted by us (\\integral/ISGRI, \\suzaku/PIN), where the 20--100 keV flux was computed from a power-law model fit (which was done in the 50--150 keV range for the OSSE data) and the uncertainty was determined from the relative error of the summed count rate in the fitted band. Fig.\\ \\ref{longest}(b) shows the 20--100 keV fluxes from the \\gro/BATSE in 7-d bins (together with those from OSSE). They have been obtained from the 20--70 keV fluxes given by \\citet{Parsons1998} by multiplying by 1.1 (assuming $\\Gamma=1.8$) and then dividing by the normalization factor with respect to OSSE of 1.42 \\citep{Parsons1998}.  Fig.\\ \\ref{longest}(b) gives the 1.5--12 keV count rate of \\xte/ASM\\footnote{http://xte.mit.edu/asmlc/ASM.html} in 30-d bins. The blue curve in Fig.\\ \\ref{longest}(c) shows the optical data from the Crimean observatory in the V band (5500 \\AA, \\citealt{Czerny2003}), which, after MJD 52000, are supplemented by the \\integral/OMC data. Since the V band data do not overlap much in time with the ASM data, we also show the 5117 \\AA\\ fluxes from various observatories given by \\citet{Shapovalova2008}.  \\begin{figure} \\includegraphics[width=\\columnwidth]{ngc4151_longest.eps} \\caption{Light curves of NGC 4151 since 1972. (a) The 20--100 keV flux from the compilation by \\citet{Perotti1991}, magenta diamonds, and from \\granat/SIGMA (cyan squares), \\gro/OSSE (red filled circles), \\sax/PDS (green triangles), \\integral/ISGRI (blue dots), and \\suzaku/PIN (the orange open circle). The dotted lines show the extrema of the 20--100 keV flux from ISGRI. (b) The 20--100 keV flux from \\gro/BATSE (green error bars) together with that from OSSE (red filled circles; same as in a). (c) The \\xte/ASM count rate in 30-d bins. (d) The 5500 \\AA\\ flux (\\citealt{Czerny2003}, blue lines), and from \\integral/OMC (vertical blue bars at MJD $> 52000$), and the 5117 \\AA\\ flux (\\citealt{Shapovalova2008}, black dots). } \\label{longest} \\end{figure}  The 1.5--12 keV count rate appears well correlated with the optical fluxes, see Fig.\\ \\ref{longest}. As found by \\citet{Czerny2003}, the medium-energy X-ray flux of NGC 4151 varies on time scales $\\sim 5$--$10^3$ d, whereas the optical variability is also present on longer time scales. For hard X-rays, a correlation with the optical and the 1.5--12 keV fluxes is less clear. The main peak of the optical emission at MJD $\\simeq 49000$--51000 (1993--1998) is reflected in the 1.5--12 keV flux but not in the 20--100 keV one (including the BATSE data). The later data from \\sax, \\integral\\/ and \\suzaku\\/ show an overall agreement with the optical and softer X-ray fluxes in a sense that the minima at MJD $\\sim$52000 and MJD $\\sim$54000 and the maximum at MJD $\\sim$52800 appear for all these bands. However, the scarce hard X-ray coverage prevents unambiguous conclusions.  \\begin{figure} \\includegraphics[width=\\columnwidth]{ngc4151_monitor.eps} \\caption{Medium and hard X-ray light curves from \\integral/ISGRI compared to those from \\xte\\/ and \\swift\\/ (shown with black circles and error bars). The 18--50 keV ISGRI flux is shown by open circles on all panels. The colours identify the three main states, bright (blue), medium (green), dim (red) and, additionally, data excluded from the spectral analysis, Revs.\\ 0076 and 0809--0811 (magenta), and the rest (cyan). (a) The \\xte/ASM count rate in 30-d bins (right axis). (b) The \\xte/PCU2 count rate (right axis). The vertical lines delineate the two periods for which the dim state PCA spectrum was obtained. (c) Comparison with the \\swift/BAT 15--50 keV flux in 7-d bins. The vertical lines show the times of the X4 (dim), X5 (medium) \\xmm\\/ observations and that of \\suzaku\\/ (dim). The horizontal dashed lines show the average flux levels of ISGRI bright, medium and dim states. } \\label{long} \\end{figure}  The results shown in Fig.\\ \\ref{longest}(a--b) show that the hard X-ray emission of NGC 4151 has varied in well defined limits over the last 39 years. Almost all fluxes are within the extreme 20--100 keV fluxes from ISGRI (with the maximum in 2003 May and the minimum in 2007 May). There are four flux measurements from older (1975--77) observations above this maximum. However, they are relatively doubtful, corresponding to marginal ($\\approx 2$--$3\\sigma$) detections at $\\ga$ 40 keV. Thus, ISGRI 20--100 keV flux range appears very close to the corresponding overall actual range. This is also supported by the BATSE data, Fig.\\ \\ref{longest}(b), varying within this range. We also notice that NGC 4151 was rarely seen at low hard X-ray fluxes before the \\integral\\/ launch.  Fig.\\ \\ref{long} compares ISGRI 18--50 keV fluxes with the contemporaneous data in the medium and hard X-ray bands from \\xte\\/ and \\swift. The ISGRI fluxes are well correlated with those from both \\xte\\/ and the \\swift/BAT\\footnote{http://swift.gsfc.nasa.gov/docs/swift/results/transients}. In particular, we see a good correlation with the 1.5--12 keV rate in spite of the variable absorption affecting a lower part of this band. Fig.\\ \\ref{long} also shows that a large fraction of the \\integral\\/ observations happened during periods with very low hard X-ray flux. Fig.\\ \\ref{long} also identifies the \\integral\\/ observations used for the spectral sets (Table \\ref{obsint}).  \\begin{figure} \\includegraphics[width=\\columnwidth]{ngc4151_short.eps} \\caption{Illustration of short time-scale variability. The blue bars (left axis) and the red dots show the 18--40 keV ISGRI flux and the optical OMC flux, respectively, during the dedicated \\integral\\/ observations. The vertical lines and the green arrows in the top panel show the times of the bright-state X1--X3 \\xmm\\/ and 11 \\xte/PCA observations, respectively. The data used for the bright \\integral\\/ spectrum are those in the top panel simultaneous with the OMC observations except those around MJD 52788 (Rev.\\ 0076), which have markedly lower fluxes. } \\label{short} \\end{figure}  The BAT light curve, Fig. \\ref{long}(c), shows that the hard X-ray flux varies by a factor of a few on a week-time scale. When the source is sufficiently bright, it is possible for ISGRI to monitor hour time-scale variability. This was the case for the bright state of 2003 May, and the source showed such variability, see Fig.\\ \\ref{short}. The shown ISGRI and OMC light curves were extracted with a time bin equal to a single pointing duration (10 m--2.5 h, typically $\\sim 1$ h). When observed in dimmer states, ISGRI flux remains constant within the measurement errors during a given continuous observation. The optical flux appears constant on hour--day time scales. Still, it varies on longer time scales, and it increases by $\\sim$30 per cent when the source is bright in X-rays (2003 May, 2009 June).  We have extracted the bright (X1--X3) \\xmm/EPIC pn light curves in several energy bands in 100-s time bins, and compared with the OMC 100-s light curves. No variability is seen in either light curves on time scales $\\la 1$ h. The \\xmm\\/ light curves vary slowly in good agreement with the trends observed for ISGRI (and also JEM-X) fluxes.  \\begin{figure} \\includegraphics[width=\\columnwidth]{ngc4151_correl.eps} \\caption{(a) The OMC 5500 \\AA\\ flux vs.\\ the ISGRI 18--40 keV flux. (b) The JEM-X 3--10 keV flux vs.\\ the ISGRI flux. Each cross corresponds to a single \\integral\\/ pointing. The colours identify the three main states, bright (blue), medium (green), dim (red), and, additionally, Revs.\\ 0076, 0809--0811 (magenta). } \\label{correl} \\end{figure}  The \\integral\\/ data are suitable for correlation studies because they provide simultaneous data in the optical, medium and hard X-ray bands. Here, we use all the observations with the off-axis angle $<3\\degr$ for which OMC and JEM-X data are available. They include also part of the observations not used in our spectral analysis, namely those identified with magenta circles in Fig.\\ \\ref{long}. Fig.\\ \\ref{correl} shows the correlations between the 18--40 keV flux and those fluxes in the 3--10 keV and V bands. Both show strong correlations, but the correlation patterns are different. The optical fluxes fall into two separate regions, within each they are correlated with the hard X-rays. Significant rank-order probabilities ($>\\! 0.999$) of the ISGRI/OMC correlation are found separately for each of the regions, but also and for the entire data set. The observed bimodality may result from strong optical long-term variability between the X-ray observations. On the other hand, the 3--10 keV and 18--40 keV fluxes exhibit a single, extremely strong, linear correlation. Statistical uncertainty blurs somewhat the correlation for low fluxes. This result indicates that during all \\integral\\/ observations the slope of the X-ray spectrum of NGC 4151 below 40 keV remains approximately constant.   \\section[]{Spectral analysis} \\label{spectral}  \\subsection{Assumptions and selection of spectral sets} \\label{assumptions}  Spectral fitting is performed with {\\sc xspec} 11.3 \\citep{Arnaud1996}. Errors are given for 90 per cent confidence level for a single parameter, $\\Delta \\chi^{2} = 2.7$. The luminosity and distance, $D=13.2$ Mpc, are for $H_{0} = 75$ km s$^{-1}$ Mpc$^{-1}$. We use the elemental abundances of \\citet{Anders1982} and the photoelectric absorption cross-sections from \\citet{Balucinska1992}. The Galactic column density is set to $N_{\\rm H}^{\\rm G}= 2.1\\times 10^{20}$ cm$^{-2}$. The assumed inclination angle of the reflector and the hot plasma (for anisotropic geometries) is $i=45\\degr$ \\citep{Das2005}. The reflector is assumed to be neutral, and relativistic broadening of reflected spectra is neglected. Still, in some cases we test for the effects of either a variable Fe abundance of the intrinsic absorber and the reflector, a variable $i$, or an ionization of the reflector.  The main assumption used in our selection of spectral sets is that the state of the emitting hot plasma is determined by the hard X-ray flux. This is justified by our finding that the shape of ISGRI spectra for a given revolution is almost constant at a given hard X-ray flux. As an additional constraint, we do not split \\integral\\/ data from a single revolution into different spectral sets. Based on this criterion, we select three average ISGRI spectra, as specified in Table \\ref{obsint} and Fig.\\ \\ref{long}. These spectra are accompanied by spectra taken with the other instruments. The state is assigned to each of them on the basis of the light curves presented in Figs.\\ \\ref{long}--\\ref{short}. In particular, the high quality of the \\swift/BAT data allows us to determine the hard X-ray flux for periods with no \\integral\\/ observations, and thus to assign the correct state to the existing lower-energy data sets (Table \\ref{obsother}). On the other hand, the data from \\xte\\/ determine by themselves the hard X-ray state.  The ISGRI bright (B) and dim (D) spectra are of relatively good quality, and correspond to the extrema of the measured hard X-ray flux. For the flux range in-between, we identify two flux ranges. However, we had no corresponding low-energy data sets for the upper of them. Thus, we use the lower of them, denoted as medium (M). The fitted data sets are as follows.  The bright state, B. We use the data from 2003 May made at low off-axis angles, as shown in the top panel of Fig.\\ \\ref{short}. The spectrum includes a part of Rev.\\ 0074 (without its beginning) and Rev.\\ 0075. For these data, we have simultaneous \\xte\\/ and \\xmm\\/ observations, see Table \\ref{obsother}. The fitted spectral set consists of ISGRI spectrum, the X1+X2 and X3 \\xmm\\/ spectra, and the \\xte\\/ PCU2 spectrum summed over 11 observations.  The medium state, M. We sum ISGRI data from Revs.\\ 0634, 0636, 0678--0679, and use the summed JEM-X 1 spectrum from Revs.\\ 0634 and 0636, and the X5 \\xmm\\/ spectrum.  The dim state, D. The summed ISGRI spectrum is from Revs.\\ 0310--0563. We use the \\xte\\/ PCU2, the X4 \\xmm\\/ EPIC pn, and the \\suzaku\\/ XIS0,3 and XIS1 spectra.  We have also investigated the effect of including other available spectra for each of the data set, whose resulting sets we denote 'extended'. The extended set B also includes the spectra from the \\integral/JEM-X 2 and SPI, and the \\xte/HEXTE clusters 0, 1. The extended set M also includes the SPI spectrum. The extended data set D also includes two \\integral\\/ JEM-X 1 summed spectra, from Revs.\\ 0521--0522 and 0561--0563, the \\suzaku\\/ PIN spectrum, the \\xte\\/ HEXTE cluster 1 spectrum, and the corresponding SPI spectrum. We discuss the effect of using the extended data sets in Section \\ref{Compton} below.  In order to limit the complexity of the fitted models, we have decided not to use the X-ray spectra at $E< 2.5$ keV. A two-component partial absorption was then sufficient to model absorption of the continuum. Using the X-ray spectra in the 0.1--2.5 keV band would have required to add a number of spectral lines as well as a soft excess component. As we have tested, including that low-energy band would not affect our determination of the parameters of the hot plasma and of Compton reflection.  Thus, the EPIC pn, XIS0+XIS3, and XIS1 spectra were used in the 2.5--11.3 keV, 2.5--9.5 keV and 2.5--9.0 keV bands, respectively. The PCA/PCU2 and JEM-X spectra were fitted in the 3--16 keV and 4--19 keV bands, respectively. The HXD/PIN spectrum was used in the 19--60 keV band because the $<$19-keV data were showing a strong excess. For ISGRI, data below 20 keV were always excluded, whereas the high-energy limit was in the 180--200 keV range, depending on our $3\\sigma$ detection threshold used to select good spectral channels. The ISGRI spectra in the staring mode from Revs.\\ 0074--75 show an excess of $\\sim$10 per cent below 23 keV, whose origin remains unknown. It could be related to a specific setting of the low-energy threshold in the early period of the mission. Thus, we used only the data at $>23$ keV in this case. The HEXTE spectra were used in the 13--160 keV band for the bright state and in 23--110 keV for the dim state. The SPI spectra were fitted in the 23--200 keV band.  While fitting a given spectral set, the parameters of intrinsic continua are the same for all included spectral data, allowing only the normalization of the model to vary between them. This reflects our assumption that the hard X-ray flux determines the intrinsic state of the source, even for non-simultaneous data. On the other hand, absorption and the Fe K emission may vary quickly \\citep{Puccetti2007}, and we cannot expect their parameters to be the same for different low-energy spectra associated with the same spectral set. Thus, we have allowed them to vary between different spectra. Our assumed absorber consists of one fully covering the source with the column density of $N_{\\rm H}^{\\rm i}$, and the ionization parameter, $\\xi \\equiv 4\\upi F_{\\rm i}/n$, where $F_{\\rm i}$ is the 5 eV--20 keV irradiating flux and $n$ is the density of the reflector. For the low-resolution detectors, PCA and JEM-X, we assumed $\\xi=0$. For the high-resolution detectors, EPIC pn and XIS, we allowed $\\xi>0$, and, in addition, we included partial-covering neutral absorption, with the column density of $N_{\\rm H}^{p}$ covering a fraction of $f^{\\rm p}$ of the flux. For the ionized absorption, we use the model by \\citet{Done1992}. Although it is not applicable to highly ionized plasma, it is sufficiently accurate at low/moderate ionization, such as that found in our data. Absorption for ISGRI, HEXTE, SPI and PIN spectra is only marginally important, and we thus used only the neutral fully-covering $N_{\\rm H}^{\\rm i}$ at the values fitted to the corresponding PCA or JEM-X spectra. The Fe K$\\alpha$ emission is modelled by a Gaussian line, with the centre energy, width, photon flux and equivalent width of $E_{\\rm Fe}$, $\\sigma_{\\rm Fe}$, $I_{\\rm Fe}$ and $w_{\\rm Fe}$, respectively. Fitting the dim state spectra, we have found that adding an Fe K$\\beta$ line (at $\\approx$7.1 keV) was significantly improving the fit for the EPIC pn and XIS spectra.  \\subsection{The cut-off power-law model} \\label{cutoff}  We first fit an e-folded power-law model including a Compton reflection component. This allows us to compare our results to other published ones. We use the {\\sc pexrav} model \\citep{Magdziarz1995}, and our other assumptions are specified above. The results are shown in Table \\ref{pl_fits}, where we do not show the fitted absorber and line parameters since they are similar to those for our main model of thermal Comptonization, see Section \\ref{Compton} below.  Our results regarding $\\Gamma$ and $R$ are similar to those of earlier works. However, a significant difference is found for the e-folding energy, $E_{\\rm c}$, which we find to be high, and compatible with no cut-off for two of our states, see Table \\ref{pl_fits}. A finite $E_{\\rm c}$ is found only for the bright state. Even then, we obtain values well above those reported before, which are 30--45 keV (P01), 80--180 keV \\citep{DeRosa2007} and $\\sim$210 keV (Z02). As discussed in \\citet{Zdziarski2003}, an exponential cut-off is much shallower than that characteristic of thermal Comptonization, and not sharp enough to model the spectral high-energy cut-offs observed in Seyferts, in particular in NGC 4151. Thus, our high values of $E_{\\rm c}$ result from the good quality of the data below the cut-off, which dominate the statistics and force an approximately straight power law to continue to just below the beginning of the cut-off.  \\begin{table} \\setlength{\\tabcolsep}{1.2mm} \\begin{center}{ \\caption{Spectral fitting results for the e-folded power-law model including reflection. The {\\sc xspec} model is {\\sc constant*wabs*absori*zpcfabs(pexrav\\ +zgauss)}, and the normalization, $K$, is given by the photon flux at 1 keV in units of $10^{-2}$ keV$^{-1}$ cm$^{-2}$ s$^{-1}$. } \\small \\label{pl_fits} \\begin{tabular}{cccccc} \\hline State & $E_{\\rm c}$ [keV] & $\\Gamma$ & $R$ & $K$ & $\\chi ^{2}$/d.o.f. \\\\ \\hline Bright & 264$_{-26}^{+48}$ & 1.71$_{-0.01}^{+0.06}$ & 0.45$_{-0.05}^{+0.12}$ & 8.7$_{-0.3}^{+1.1}$ & 3179/3414 \\\\ Medium & $>$ 1025 & 1.81$_{-0.03}^{+0.05}$ & 0.6$_{-0.3}^{+0.3}$ & 4.6$_{-0.7}^{+1.4}$ & 1554/1599 \\\\ Dim & $>$ 1325 & 1.81$_{-0.01}^{+0.01}$ & 0.92$_{-0.05}^{+0.05}$ & 2.24$_{-0.02}^{+0.01}$ & 3338/3384 \\\\ \\hline \\end{tabular} }\\end{center} \\end{table}  \\subsection{Comptonization model} \\label{Compton}  \\begin{table*} \\begin{minipage}{172mm} \\begin{center} \\caption{{\\bf (a)} Spectral fitting results with thermal Comptonization as the main continuum. The {\\sc xspec} model is {\\sc constant*wabs*absori*zpcfabs*(compps\\ +zgauss)}. See Section \\ref{Compton} for the definitions of the parameters. $L_{\\rm X\\gamma}$ and $L_{\\rm X\\gamma}'$ are the bolometric model luminosity (in units of $10^{43}$ erg s$^{-1}$) with and without the reflection component, respectively, and only the values for the geometry 0 are given as the values for other geometries are similar. } \\small \\label{comp_results} \\begin{tabular}{lccccccccc} \\hline State & $kT_{\\rm e}$ [keV] & $y$ & $\\tau$ & $R$ & $K\\,[10^8]$ & $L_{\\rm X\\gamma}$ & $L_{\\rm X\\gamma}'$ &Geometry & $\\chi ^{2}$/d.o.f. \\\\ \\hline Bright & 54$_{-5}^{+7}$ & 1.10$_{-0.02}^{+0.02}$ & 2.6$_{-0.3}^{+0.3}$ & 0.40$_{-0.06}^{+0.07}$ & 5.2$_{-0.6}^{+1.4}$ & 5.18 & 4.61 & sphere (0) & 3165/3414 \\\\ & 62$_{-7}^{+7}$ & 0.60$_{-0.01}^{+0.01}$ & 1.3$_{-0.1}^{+0.1}$ & 0.55$_{-0.07}^{+0.08}$ & 7.3$_{-0.9}^{+1.9}$ & & & slab (1) &  3166/3414 \\\\ & 62$_{-7}^{+9}$ & 0.61$_{-0.01}^{+0.01}$ & 1.3$_{-0.2}^{+0.1}$ & 0.55$_{-0.06}^{+0.09}$ & 3.9$_{-0.7}^{+0.4}$ & & & slab ($-1$) & 3166/3414 \\\\ & 73$_{-9}^{+14}$ & 1.06$_{-0.02}^{+0.02}$ & 1.9$_{-0.2}^{+0.4}$ & 0.47$_{-0.07}^{+0.05}$ & 12.0$_{-1.6}^{+2.9}$ & & & cylinder (2) & 3168/3414 \\\\ & 73$_{-13}^{+13}$ & 1.06$_{-0.01}^{+0.03}$ & 1.9$_{-0.3}^{+0.3}$ & 0.51$_{-0.9}^{+0.5}$ & 5.4$_{-1.1}^{+1.4}$ & & & cylinder ($-2$) & 3168/3414 \\\\ & 73$_{-7}^{+16}$ & 1.21$_{-0.02}^{+0.02}$ & 2.1$_{-0.6}^{+0.4}$ & 0.48$_{-0.06}^{+0.06}$ & 11.9$_{-1.4}^{+3.2}$ & & & hemisphere (3) & 3168/3414 \\\\ & 73$_{-9}^{+15}$ & 1.21$_{-0.02}^{+0.03}$ & 2.1$_{-0.4}^{+0.3}$ & 0.50$_{-0.06}^{+0.08}$ & 4.7$_{-1.0}^{+1.5}$ & & & hemisphere ($-3$) & 3168/3414 \\\\ & 61$_{-5}^{+7}$ & 0.80$_{-0.02}^{+0.01}$ & 1.7$_{-0.2}^{+0.1}$ & 0.38$_{-0.05}^{+0.04}$ & 1.2$_{-0.1}^{+0.1}$ & & & sphere (4) & 3172/3414 \\\\ & 57$_{-7}^{+13}$ & 0.86$_{-0.01}^{+0.01}$ & 1.9$_{-0.4}^{+0.2}$ & 0.39$_{-0.05}^{+0.06}$ & 6.4$_{-0.6}^{+0.6}$ & & & sphere ($-4$) & 3169/3414 \\\\ & 57$_{-5}^{+10}$ & 0.86$_{-0.01}^{+0.01}$ & 1.9$_{-0.3}^{+0.2}$ & 0.38$_{-0.05}^{+0.06}$ & 4.9$_{-0.4}^{+1.5}$ & & & sphere ($-5$) & 3166/3414 \\\\ \\hline Medium & 128$_{-24}^{+74}$ & 1.02$_{-0.13}^{+0.08}$ & $1.0_{-0.3}^{+0.5}$ & $< 0.41$ & 3.9$_{-1.0}^{+0.4}$ &2.26 & 2.16 & sphere (0) & 1554/1597 \\\\ \\hline Dim & 190$_{-12}^{+13}$ & 0.98$_{-0.01}^{+0.01}$ & 0.66$_{-0.05}^{+0.04}$ & 0.75$_{-0.04}^{+0.04}$ & 1.54$_{-0.01}^{+0.01}$ &1.47 & 1.19 & sphere (0) & 3337/3384 \\\\ & 200$_{-8}^{+7}$ & 0.49$_{-0.01}^{+0.01}$ & 0.31$_{-0.01}^{+0.01}$ & 1.01$_{-0.05}^{+0.05}$ & 2.04$_{-0.02}^{+0.01}$ & & & slab (1) & 3335/3384 \\\\ & 188$_{-7}^{+6}$ & 0.48$_{-0.01}^{+0.01}$ & 0.33$_{-0.01}^{+0.01}$ & 1.04$_{-0.05}^{+0.05}$ & 1.20$_{-0.01}^{+0.01}$ & & & slab ($-1$) & 3350/3384 \\\\ & 210$_{-14}^{+12}$ & 1.05$_{-0.01}^{+0.01}$ & 0.64$_{-0.04}^{+0.04}$ & 0.79$_{-0.04}^{+0.04}$ & 2.85$_{-0.01}^{+0.02}$ & & & cylinder (2) & 3334/3384 \\\\ & 215$_{-13}^{+12}$ & 1.01$_{-0.01}^{+0.01}$ & 0.60$_{-0.03}^{+0.04}$ & 0.86$_{-0.05}^{+0.05}$ & 1.73$_{-0.01}^{+0.01}$ & & & cylinder ($-2$) & 3336/3384 \\\\ & 227$_{-14}^{+10}$ & 1.21$_{-0.02}^{+0.01}$ & 0.68$_{-0.03}^{+0.04}$ & 0.84$_{-0.04}^{+0.04}$ & 2.89$_{-0.02}^{+0.01}$ & & & hemisphere (3) & 3334/3384 \\\\ & 220$_{-11}^{+13}$ & 1.14$_{-0.01}^{+0.01}$ & 0.66$_{-0.04}^{+0.03}$ & 0.90$_{-0.05}^{+0.05}$ & 1.71$_{-0.01}^{+0.02}$ & & & hemisphere ($-3$) & 3336/3384 \\\\ & 186$_{-9}^{+11}$ & 0.66$_{-0.01}^{+0.01}$ & 0.45$_{-0.03}^{+0.02}$ & 0.77$_{-0.04}^{+0.04}$ & 0.866$_{-0.004}^{+0.004}$ & & & sphere (4) & 3341/3384 \\\\ & 191$_{-17}^{+13}$ & 0.86$_{-0.01}^{+0.01}$ & 0.58$_{-0.04}^{+0.05}$ & 0.77$_{-0.04}^{+0.04}$ & 1.56$_{-0.01}^{+0.01}$ & & & sphere ($-4$) & 3337/3384 \\\\ & 196$_{-13}^{+13}$ & 0.81$_{-0.01}^{+0.01}$ & 0.53$_{-0.04}^{+0.04}$ & 0.77$_{-0.04}^{+0.04}$ & 1.49$_{-0.01}^{+0.01}$ & & & sphere ($-5$) & 3337/3384 \\\\ \\hline \\end{tabular} \\end{center} \\end{minipage} \\end{table*}  \\setcounter{table}{3} \\begin{table*} \\begin{minipage}{172mm} \\begin{center} \\caption{{\\bf (b)} Continuation of Table \\ref{comp_results}(a). Fit results with the thermal Comptonization model (for a sphere, the geometry parameter $=0$) regarding the absorber, line component, and the relative normalization, $C_{\\rm ISGRI}$, of a given X-ray spectrum with respect to ISGRI. See Section \\ref{assumptions} for the definitions of the other parameters. The units of $N_{\\rm H}$, $\\xi$, and $I_{\\rm Fe}$ are $10^{22}$ cm$^{-2}$, $10^{-2}$ erg cm s$^{-1}$, $10^{-4}$ cm$^{-2}$ s$^{-1}$, respectively, and `f' denotes a fixed parameter. } \\begin{tabular}{llccccccccc} \\hline State & X-ray spectrum & $C_{\\rm ISGRI}$ & $N_{\\rm H}^{\\rm i}$ & $\\xi$ & $N_{\\rm H}^{\\rm p}$ & $f^{\\rm p}$ & $E_{\\rm Fe}$ [keV] & $\\sigma_{\\rm Fe}$ [keV] & $I_{\\rm Fe}$ & $w_{\\rm Fe}$ [eV] \\\\ \\hline Bright & EPIC pn (X1+X2) & 0.92$_{-0.03}^{+0.03}$ & 4.9$_{-0.7}^{+0.9}$ & 10$_{-6}^{+79}$ & 15.8$_{-3.3}^{+3.9}$ & 0.39$_{-0.07}^{+0.09}$ & 6.400$_{-0.011}^{+0.011}$ & 0.07$_{-0.01}^{+0.02}$ & 3.0$_{-0.3}^{+0.3}$ & { 83$_{{{-9}}}^{{{+7}}}$ } \\\\ & EPIC pn (X3)    & 1.08$_{-0.03}^{+0.03}$ & 3.0$_{-1.8}^{+1.5}$ & 64$_{-61}^{+47}$ & 9.5$_{-1.1}^{+4.0}$ & 0.55$_{-0.22}^{+0.18}$ & 6.393$_{-0.012}^{+0.013}$ & 0.07$_{-0.01}^{+0.01}$ & 2.8$_{-0.3}^{+0.3}$ & { 76$_{{{-9}}}^{{{+9}}}$ } \\\\ & PCA PCU2 (B)    & 1.17$_{-0.03}^{+0.03}$ & 7.4$_{-0.4}^{+0.4}$ & 0f & 0f & -- & 6.06$_{-0.09}^{+0.08}$ & 0.36$_{-0.13}^{+0.15}$ & 6.8$_{-0.8}^{+1.3}$ & { 169$_{{{-20}}}^{{{+32}}}$ } \\\\ \\hline Medium & EPIC pn (X5) & 0.83$_{-0.06}^{+0.08}$ & 7.6$_{-0.4}^{+0.5}$ & 4$_{-3}^{+11}$ & 22.6$_{-1.5}^{+4.4}$ & 0.50$_{-0.04}^{+0.05}$ & 6.388$_{-0.007}^{+0.007}$ & 0.06$_{-0.01}^{+0.01}$ &  2.7$_{-0.3}^{+0.5}$ & { 149$_{{{-12}}}^{{{+27}}}$ } \\\\ & JEM-X (M)    & 0.74$_{-0.06}^{+0.06}$ & 13.0$_{-3.3}^{+3.7}$ & 0f & 0f & -- & -- & -- & -- & -- \\\\ \\hline Dim & EPIC pn (X4) & 1.17$_{-0.01}^{+0.01}$ & 6.36$_{-0.15}^{+0.14}$ & 4$_{-3}^{+30}$ & 23.5$_{-0.7}^{+0.6}$ & 0.69$_{-0.01}^{+0.01}$ & 6.398$_{-0.005}^{+0.004}$ & 0.06$_{-0.01}^{+0.01}$ & 2.3$_{-0.1}^{+0.1}$ & { 271$_{{{-13}}}^{{{+13}}}$ }  \\\\ & XIS0,3 (D)   & 1.11$_{-0.01}^{+0.01}$ & 8.03$_{-0.16}^{+0.12}$ & 0.7$_{-0.4}^{+0.9}$ & 28.8$_{-1.0}^{+0.9}$ & 0.50$_{-0.01}^{+0.01}$ & 6.384$_{-0.003}^{+0.003}$ & 0.03$_{-0.01}^{+0.01}$ & 2.2$_{-0.1}^{+0.1}$ & { 291$_{{{-13}}}^{{{+13}}}$ } \\\\ & XIS1 (D)    & 1.13$_{-0.01}^{+0.01}$ & \\multicolumn{4}{c}{fitted together with XIS0,3} & 6.412$_{-0.004}^{+0.005}$ & 0.04$_{-0.01}^{+0.01}$ & 2.3$_{-0.1}^{+0.1}$ & { 280$_{{{-13}}}^{{{+13}}}$ } \\\\ & PCA PCU2 (D) & 1.16$_{-0.01}^{+0.01}$ &  13.6$_{-0.5}^{+0.5}$ & 0f & 0f & -- & 6.4f & 0f & 2.1$_{-0.3}^{+0.3}$ & { 303$_{{{-39}}}^{{{+39}}}$ } \\\\ \\hline \\end{tabular} \\end{center} \\end{minipage} \\end{table*}  To model thermal Comptonization, we use the {\\sc compps} model \\citep{Poutanen1996} in {\\sc xspec}. It models Compton scattering in a plasma cloud of given temperature, $T_{\\rm e}$, and Thomson optical depth, $\\tau$, for a number of geometries and locations of the seed photon sources. Since $T_{\\rm e}$ and $\\tau$ are strongly intrinsically anticorrelated, we also use the Compton parameter (e.g., \\citealt{Rybicki1979}), $y\\equiv 4(kT_{\\rm e}/m_{\\rm e} c^2)\\tau$ (where $m_{\\rm e}$ is the electron mass) as the second fitting parameter instead of $\\tau$. The advantage of this choice is that $T_{\\rm e}$ and $y$ are almost orthogonal; $y$ determines closely the power-law slope at low energies whereas $kT_{\\rm e}$ determines the position of the high-energy cut-off. The seed photons are assumed to have a disc blackbody distribution ({\\sc diskbb} in {\\sc xspec}, \\citealt{Mitsuda84}) with the maximum temperature of $kT_{\\rm bb}=10$ eV. The model normalization, $K$, is that of the seed disc blackbody model as defined in {\\sc xspec}, see Section \\ref{xray}.  Among the geometries included in {\\sc compps}, we consider here four spherical cases; an approximate treatment of radiative transfer using escape probability (which is denoted in {\\sc compps} by the geometry parameter $=0$), a sphere with central soft photons (geometry 4), homogeneously distributed seed photons (geometry $-4$) and seed photons distributed according to the diffusion-equation eigen-function, $\\propto \\sin(\\upi\\tau'/\\tau)/ (\\upi\\tau'/\\tau)$, where $0\\leq \\tau'\\leq\\tau$ ($-5$). Then, we consider a slab with the seed photons either at its bottom (geometry 1) or distributed homogeneously (geometry $-1$). Also, we consider the hot plasma in the shape of either a cylinder with the height equal to its radius, or a hemisphere, with the seed photons being either at its bottom (2, 3, respectively) or homogeneously distributed ($-2$, $-3$, respectively).  The spectral data of NGC 4151 studied in this paper are of high quality and represent the largest set ever collected for this source. This gives us a possibility to test how precise information about X$\\gamma$ emission from brightest Seyferts can be achieved with the current satellites. Therefore, we extensively test various variants of the Comptonization model for our two best-quality states, bright and dim. This could, in principle, give us some indications regarding the actual geometry of the source. Our results are presented in Tables \\ref{comp_results}(a) and (b), regarding the parameters of the Comptonizing plasma and the strength of reflection (which are the same for all used detectors within a given state), and of the absorber and the Fe line (which are specific to a given X-ray detector within each of the states), respectively. The results in Table \\ref{comp_results}(b) and the entries for the bolometric model luminosity, $L_{\\rm X\\gamma}$, in Table \\ref{comp_results}(a) are given only for the case of spherical geometry calculated using escape probability formalism (geometry 0).  We find the thermal Comptonization model to provide very good fits to the data. The values of the reduced $\\chi^2$ are $\\la 1$, which, however, appears not to be due to an overly complex model. In particular, the high-energy continua are determined only by three parameters, $y$, $R$, $kT_{\\rm e}$. For the bright state, Comptonization provides a much better fit than the e-folded power law, with $\\Delta \\chi^2$ of up to $-14$. However, any trends seen for the geometry of the Comptonizing plasma are rather weak. In the case of the bright state, equally good models can be obtained with either a slab or a sphere. Also, models with seed photons being distributed within the source are of similar quality as those with localized seed photons. On the other hand, the differences between the models are slightly stronger in the case of the dim state. In particular, the slab model with seed photons at its bottom is somewhat better than that with the seed photons distributed throughout it. This can be a hint that the actual source has the seed photons external to the source rather than internal (which would be the case for the dominant synchrotron seed photons). We also note that the fitted values of $kT_{\\rm e}$ are rather insensitive to the assumed geometry in both the bright and dim states. Our results for the bright state are similar to those of previous Comptonization fits, in particular to those of Z02, who obtained $kT_{\\rm e}= 73^{+34}_{-29}$ keV, $y=0.88^{+0.12}_{-0.11}$, and $R=0.60^{+0.24}_{-0.21}$ at $L_{\\rm X\\gamma}\\simeq 3.9\\times 10^{43}$ erg s$^{-1}$ using a broad-band X-ray spectrum from \\asca\\/ and OSSE.  In all geometries, we see in Table \\ref{comp_results}(a) that the Compton $y$ parameter remains approximately constant between the bright and dim states, with only a slight decrease in the dim state. For the medium state, Table \\ref{comp_results}(a) gives the results only for the model with escape probabilities, for which the best fit $y$ is almost the same as that in the dim state. These results are illustrated by the confidence contours shown in Fig.\\ \\ref{contours}.  \\begin{figure} \\includegraphics[width=\\columnwidth]{ngc4151_contours.eps} \\caption{The 90 per cent confidence regions for the Compton $y$ parameter (left) and the optical depth, $\\tau$ (right), vs.\\ the temperature, $kT_{\\rm e}$ for three geometries of the Comptonizing plasma. The solid blue, solid red and green dashed (only for $y$) contours correspond to the bright, dim and medium state, respectively. The dashed curves for $\\tau$ vs.\\ $kT_{\\rm e}$ correspond to $\\tau\\propto kT_{\\rm e}^{-1}$, normalized to the bright-state results.} \\label{contours} \\end{figure}  Apart from the results given in Table \\ref{comp_results}, a reduction of $\\chi^{2}$ appears when the Fe abundance (assumed equal for the absorber and the reflector) is a free parameter. We found a moderate overabundance, ${1.5\\pm 0.2}$ times that of \\citet{Anders1982} for the bright state, with $\\Delta\\chi^{2} = -6$, and ${1.3\\pm 0.1}$ for the dim state, $\\Delta \\chi^{2} = -2$. For the \\xmm\\/ observation in 2000 December, an Fe overabundance by 2--3 was claimed by \\citet{Schurch2003} but that fit was done with the \\xmm\\/ spectrum only and the reflection strength (affecting that determination) was not well constrained.  We also find that both the bright and dim state spectra prefer an inclination angle $>45\\degr$, with $\\Delta \\chi^{2} = -3$ and $-6$, respectively, at $75\\degr$. Since this effect is observed for symmetric (sphere) and asymmetric (hemisphere) geometries, we conclude that its main cause is the changing shape of the reflection component. On the other hand, allowing the reflector to be ionized does not improve the fit, and $\\xi < 0.2$ erg cm s$^{-1}$ (at an assumed reflector temperature of $10^5$ K).  The lack of the far-UV and very soft X-ray spectra does not allow us to fit the maximum temperature of the seed disc blackbody photons, $kT_{\\rm bb}$. We have found that changing it to either 5 or 20 eV yields no improvement of the fit. This is indeed expected given that the photon energies emitted by the disc are much below the fitted energy range, at which the shape of the Comptonization spectrum has already achieved a power-law shape independent of $kT_{\\rm bb}$.  Table \\ref{comp_results}(b) also shows the normalization of a given X-ray spectrum relative to that from ISGRI. For the bright state, when the \\integral, \\xmm\\/ and \\xte\\/ observations were almost contemporary, $C_{\\rm ISGRI}$ is very close to 1 and, for both of the \\xmm\\/ spectra, it follows the flux changes seen in the top panel of Fig.\\ \\ref{short}. This shows that ISGRI, EPIC pn and PCA detectors are well cross-calibrated. In the case of the two other spectral sets, the observations are not simultaneous. Nevertheless, $C_{\\rm ISGRI}$ varies only within the range of 0.74--1.17, confirming the validity of our selection of the spectra for a given flux state.  \\begin{figure} \\includegraphics[width=\\columnwidth]{ngc4151_levels.eps} \\caption{The Comptonization model spectra (solid curves; the geometry parameter $=0$) of the three states of the NGC 4151, shown together with ISGRI spectra. The bright, medium and dim spectra are shown by the blue crosses, green circles and red triangles, respectively. The dashed curves show the reflection model components for the bright (upper) and dim state. The fitted JEM-X, \\xte, \\xmm\\/ and \\suzaku\\/ spectra are not shown for clarity.} \\label{models} \\end{figure}  \\begin{figure} \\includegraphics[width=\\columnwidth]{ngc4151_highenergy.eps} \\caption{The Comptonization models for the extreme states (solid curves, same as in Fig.\\ \\ref{models}) compared to the average OSSE spectrum (crosses and two $2\\sigma$ upper limits). The rightmost upper limit is the COMPTEL upper limit in the 0.75--1 MeV band \\citep{Maisack1995}. The histograms give the \\integral/PICsIT $3\\sigma$ upper limits for 1 Ms (upper) and 10 Ms (lower) effective exposure time. The PICsIT 1.34 Ms average 277--461 keV flux is shown by error bars with a diamond. The SPI spectrum obtained by \\citet{Bouchet2008} is shown by error bars with circles. Note that the 200--600 keV data point may be spurious. } \\label{high} \\end{figure}  We have also tested our results using the extended data sets (see Section \\ref{assumptions}). However, the quality of the fits and the plasma parameter uncertainty are not improved when they are used. This confirms that our selection of the primary sets was valid. Thus, we do not present here those results.  We note that even for such a bright AGN as NGC 4151, \\textbf{we have no data at energies $\\ga 200$ keV, as shown in Fig.\\ \\ref{models}. Still, for the dim state, we have found $kT_{\\rm e}\\simeq 200$ keV with a statistical uncertainty of $\\sim 7$ per cent or so. The physical reason for this low uncertainty is the observed lack of spectral curvature. In the process of thermal Comptonization, a power law slope at $E\\la kT_{\\rm e}$ appears due to the merging of many individual up-scattering profiles, each being curved. The cut-off at $E\\ga kT_{\\rm e}$ is due to the electron energies being comparable to the photon ones, in which case a photon no longer increases its energy in a scattering. Since we see no cut-off up to $\\sim$200 keV, models with a lower $kT_{\\rm e}$ are ruled out. On the other hand, models with a higher $kT_{\\rm e}$ correspond to an increase of the photon energy in a single scattering by a factor so large that the individual, curved, scattering profiles become clearly visible and the spectrum at $E\\la kT_{\\rm e}$ is no longer a power law. Since the observed spectral shape {\\it is\\/} a power law, this case is also excluded. We note, however, that our assumed model corresponds to a single scattering zone. A power law spectrum at $E\\la 200$ keV may also be produced by a superposition of emission regions with different temperatures (smoothing out the resulting spectrum), as, e.g., in a hot accretion flow. In this case, the maximum temperature of the flow could be higher than 200 keV. In any case, we do see a very clear difference in the characteristic Comptonization temperature between the bright and dim state.}  We comment here on the very high value of $kT_{\\rm e}=315\\pm 15$ keV obtained by P01 for the 1999 January \\sax\\/ spectrum of NGC 4151. The fit of that model appears rather poor when compared with the curvature of the \\sax/PDS data above 50 keV, see fig.\\ 2(b) in P01. We also find that the PDS spectrum is very similar to the average OSSE spectrum \\citep{Johnson1997}, shown in Fig.\\ \\ref{high} (with only the normalization of the PDS spectrum being about 5 per cent lower). This indicates that the source was in a moderately bright state in 1999 January, for which state the temperature of 315 keV appears highly unlikely. We cannot explain this discrepancy; it may be that $kT_{\\rm e}=315\\pm 15$ keV represented a spurious local minimum.  As we see in Fig.\\ \\ref{high}, our bright-state model spectrum agrees well in shape with the average OSSE spectrum \\citep{Johnson1997}. The latter also agrees with the first three points of the average NGC 4151 spectrum from \\integral/SPI \\citep{Bouchet2008}. However, the SPI 200--600 keV flux is implausibly high, which may reflect a problem with the background estimate. Indeed, the SPI analysis of \\citet{Petry2009} yields only upper limits at $E>150$ keV.  We see in Fig.\\ \\ref{high} that the Comptonization models of the bright and dim states cross at $\\simeq 300$ keV. Unfortunately, it is impossible to verify it with our data. In particular, we cannot test if there is any non-thermal component at high energy. Even the average spectrum from OSSE, the most sensitive detector up to date at $E\\ga 150$ keV, gives only a weak detection of NGC 4151 at $E\\ga 250$ keV with an $>4$ Ms exposure time. From the shown upper limits (including systematic errors) of the \\integral/PICsIT, the current detector most sensitive at $E\\ga 250$ keV \\citep{Lubinski2009}, we see it would need $>10$ Ms exposure time to detect NGC 4151 in the 300--400 keV range. Merging all the suitable \\integral\\/ data (with the effective exposure time of 1.34 Ms taken without the bright state, for which the observation was in the staring mode, not suitable for the PICsIT), we extracted the PICsIT flux in the 277--461 keV band. Although it is not a detection (and the shown error bar does not include systematic errors), it provides a hint that the average NGC 4151 emission in this energy range is weak.  \\section[]{Discussion}  \\subsection{Reflection and Fe line} \\label{reflector}  \\subsubsection{Comparison with previous work} \\label{comparison}  A reflection component in the X-ray spectra of NGC 4151 was first found by \\citet{Zdziarski1996a}, using contemporaneous data from \\gro\\/ and \\ginga, yielding $R\\simeq 0.43$ (using a Comptonization model and assuming $i = 65\\degr$). Then, Z02 found $R\\simeq 0.6$ (at $i=17\\degr$) using \\gro\\/ and \\asca\\/ data, and P01 found $R\\simeq 0.2$--1.8 using \\sax\\/ data and a reflection model averaged over the viewing angle. A previous analysis of \\integral\\/ bright-state data from 2003 May yielded a larger reflection, $R = 0.72\\pm 0.14$ \\citep{Beckmann2005}, than that determined by us (their value of $kT_{\\rm e} = 94_{-10}^{+4}$ keV is also different from that found by us for the same data). By fitting the ISGRI and JEM-X bright-state spectra alone, we find results very similar to those with the \\xmme and \\xtee spectra included (Table \\ref{comp_results}). Therefore the differences appear to be a consequence of the substantial change of the \\integral\\/ calibration since 2005. Furthermore, \\citet{Beckmann2005} used the SPI spectra obtained in the staring observation mode, which can be affected by a large background uncertainty. Then, \\citet{DeRosa2007} found $R\\simeq 0.4$ for several \\sax\\/ spectra and $R\\simeq 0$ for the remaining ones, but using an e-folded power-law model rather than Comptonization. Then, \\citet{Schurch2002} found $R\\simeq 1.9$ using \\xmm\\/ data assuming $i=65\\degr$, but this result was obtained using the 2.5--12 keV band only and assuming a fixed $\\Gamma = 1.65$.  Regarding Fe K$\\alpha$ line, \\citet{Zdziarski1996a} found a narrow line with $I_{\\rm Fe}\\simeq (0.6$--$1.0)\\times 10^{-4}$ cm$^{-2}$ s$^{-1}$ using \\ginga\\/ data, and Z02 found a narrow line at $I_{\\rm Fe}\\simeq 2.5\\times 10^{-4}$ cm$^{-2}$ s$^{-1}$ accompanied by a weaker broad relativistic line using \\asca\\/ data. The broad component has not been observed later, whereas the narrow line flux measured by \\sax\\/ \\citep{DeRosa2007}, \\chandra\\/ \\citep{Ogle2000} and \\xmm\\/ \\citep{Schurch2003} was found to be in the range (1.3--$4.2)\\times 10^{-4}$ cm$^{-2}$ s$^{-1}$. As seen in Tables \\ref{comp_results}(a--b), our results on $R$ and $I_{\\rm Fe}$ approximately agree with the previous measurements.  \\subsubsection{Close and distant reflector} \\label{close_distant}  We see in Table \\ref{comp_results}(a) that the relative strength of reflection, $R$, is close to being twice higher in the dim state than in the bright one. (In the medium state, the amount of data is very limited, and the apparent lack of detectable reflection may be not typical to that flux level.) The increase of $R$ with decreasing flux could be a real change of the solid angle subtended by the reflector between the states. On the other hand, it could be that it is due to the contribution of a distant reflector, in particular a molecular torus, as postulated in the AGN unified model. Reflection from distant media is indeed commonly seen in Seyferts 2, e.g., \\citet{Reynolds1994}. If we attribute the increase of $R$ in the dim state due to that component, we can calculate the fractional reflection strength from the close reflector, presumably an accretion disc, and the flux reflected from the torus, which, due to its large size, is averaged over a long time scale, and thus assumed to be the same in either state.  The hypothesis of the stable geometry of the source is supported by the approximate constancy of the Compton $y$ parameter, see Table \\ref{comp_results}(a) and Fig.\\ \\ref{contours}, and of the photon index, $\\Gamma$, see Table \\ref{pl_fits}. Those parameters are closely related to the amplification factor of Comptonization (e.g., \\citealt{Beloborodov1999}), whose constancy can be most readily achieved if the system geometry is constant (e.g., \\citealt{Haardt1991,Zdziarski1999}). Thus, we consider an approximately constant relative reflection strength from the disc very likely. Observationally, we do see a strong $R$-$\\Gamma$ correlation in Seyferts and X-ray binaries, which can be interpreted by a geometrical feedback model \\citep{Zdziarski1999}, and which would require a constant disc $R$ for a given $\\Gamma$.  The unabsorbed flux in a given state, $F_{\\rm S}$, where ${\\rm S}=$ either D (dim) or B (bright), is the sum of the (unabsorbed) incident (i) and reflected (r) fluxes, \\begin{equation} F_{\\rm S}=F_{\\rm i,S}+F_{\\rm r,S}, \\label{decomposition} \\end{equation} where $F_{\\rm S}$ and $F_{\\rm r,S}$ can be calculated from the results of our spectral fitting in Section \\ref{Compton}. We assume that the observed (unabsorbed) reflected flux in each state is the sum of the close disc (d) state-dependent component and the constant distant torus (t) component, \\begin{equation} F_{\\rm r,S} = F_{\\rm d,S} + F_{\\rm t}. \\label{sum} \\end{equation} Then, assuming the constancy of the solid angle subtended by the disc reflector, or, equivalently, its reflection strength, $R_{\\rm d}$, the disc-reflected fluxes are, \\begin{equation} F_{\\rm d,S} =F_{\\rm r,S} (R_{\\rm d}/ R_{\\rm S}), \\label{disc_fluxes} \\end{equation} where $R_{\\rm S}$ is the reflection strength for a given state. We have four equations (\\ref{sum}--\\ref{disc_fluxes}) for four unknowns, $F_{\\rm d,B}$, $F_{\\rm d,D}$, $R_{\\rm d}$ and $F_{\\rm t}$, which can be readily solved. We note that they can be formulated without involving $F_{\\rm i,S}$, which implies that the solution does not depend on the choice of the energy interval in which the fluxes are measured as long as it includes the entire reflected spectrum. The solutions for $R_d$ and $F_{\\rm t}$ are, \\begin{equation} R_{\\rm d} = \\frac{(F_{\\rm r,B}-F_{\\rm r,D})R_{\\rm B}R_{\\rm D}}{F_{\\rm r,B} R_{\\rm D}-F_{\\rm r,D} R_{\\rm B}}, \\qquad F_{\\rm t} = \\frac{(R_{\\rm D}-R_{\\rm B})F_{\\rm r,B}F_{\\rm r,D} } {F_{\\rm r,B} R_{\\rm D}-F_{\\rm r,D} R_{\\rm B}}. \\label{solution} \\end{equation} Then, the fractional reflection strength of the torus can be approximately estimated as, \\begin{equation} R_{\\rm t}\\simeq  \\langle R\\rangle (F_{\\rm t}/ \\langle F_{\\rm r}\\rangle), \\label{torus} \\end{equation} where $\\langle R\\rangle$ and $\\langle F_{\\rm r}\\rangle$ are the average observed reflection strength and the average reflected flux, respectively.  Using the values obtained for thermal Comptonization using escape probability formalism (geometry parameter $=0$), we obtain $R_{\\rm d} = 0.27\\pm 0.07$, $F_{\\rm t}$ = { (8.2$\\pm$0.6)}$\\times 10^{-11}$ erg cm$^{-2}$ s$^{-1}$, and $R_{\\rm t}$ = { 0.24$\\pm$0.04} using either arithmetic or geometric averages in equation (\\ref{torus}). Similar values are obtained for other Comptonization geometries considered in Section \\ref{Compton}. We can check the robustness of our estimates by allowing the reflection strength to depend on the accretion rate, or flux. If we assume that $R_{\\rm d}$ in the dim state is a half of that in the bright state [which requires an appropriate change of equation (\\ref{disc_fluxes})], we obtain relatively similar values of the disc reflection in the bright state of $R_{\\rm d,B}$ = { 0.23$\\pm$0.12}, $F_{\\rm t}$ = { (1.1$\\pm$0.1)}$\\times 10^{-10}$ erg cm$^{-2}$ s$^{-1}$, and $R_{\\rm t}$ = { 0.32$\\pm$0.04}.  Similar reasoning can be applied to the Fe K$\\alpha$ line. The observed line photon flux in the state S, $I_{\\rm S}$, is assumed to be the sum of contributions from the disc, $I_{\\rm d,S}$ and the torus, $I_{\\rm t}$. The latter is assumed constant, and the former proportional to the ionizing flux incident on the disc. As a simplification, we assume the ionizing flux ($\\geq 7.1$ keV for neutral Fe) to be proportional to the differential continuum photon flux, $F(E)/E$, at the line centroid energy, which we denote $N_{\\rm S}$, and which is equal to $I_{\\rm S}/w_{\\rm S}$, where $w_{\\rm S}$ is the observed equivalent width in a given state. This then implies a constant disc line equivalent width, $w_{\\rm d}$. Above, for the sake of simplicity, we have dropped the indices 'Fe', used in Section \\ref{Compton} and Table \\ref{comp_results}(b). The equations are, \\begin{equation} I_{\\rm S} = I_{\\rm d,S}+I_{\\rm t}, \\qquad I_{\\rm d,S} = w_{\\rm d} N_{\\rm S}, \\label{line} \\end{equation} which can be solved for \\begin{equation} w_{\\rm d} = {I_{\\rm B}-I_{\\rm D}\\over N_{\\rm B}-N_{\\rm D}}, \\qquad I_{\\rm t} = {I_{\\rm D} N_{\\rm B}-I_{\\rm B}N_{\\rm D}\\over N_{\\rm B}-N_{\\rm D}}. \\label{line_solution} \\end{equation} The equivalent width of the torus line flux with respect to the average incident photon flux, $ \\langle N\\rangle$, can be estimated as, \\begin{equation} w_{\\rm t}={I_{\\rm t}/ \\langle N\\rangle}. \\label{wt} \\end{equation}  For the numerical values, we use here averages of the fit results for both EPIC pn spectra in the bright state, and for the EPIC pn and XIS spectra in the dim state, see Table \\ref{comp_results}(b). We find $I_{\\rm B}$ = { (2.9$\\pm$0.2)} $\\times 10^{-4}$ cm$^{-2}$ s$^{-1}$, $I_{\\rm D}$ = { (2.3$\\pm$0.1)}$\\times 10^{-4}$ cm$^{-2}$ s$^{-1}$, $N_{\\rm B}$ = { (3.6$\\pm$0.1)}$\\times 10^{-3}$ keV$^{-1}$ cm$^{-2}$ s$^{-1}$ and $N_{\\rm D}$ = { (8.1$\\pm$0.1)}$\\times 10^{-4}$ keV$^{-1}$ cm$^{-2}$ s$^{-1}$. This yields $w_{\\rm d}$ = { 23$\\pm$11} eV, $I_{\\rm t}$ = { (2.1$\\pm$0.9)}$\\times 10^{-4}$ cm$^{-2}$ s$^{-1}$, and $w_{\\rm t}$ = { 93$\\pm$41} eV, { 120$\\pm$53} eV using the arithmetic or geometric average, respectively.  The equivalent width of the Fe line at $i=45\\degr$ by an isotropically illuminated cold disc for $\\Gamma = 1.75$ is $\\simeq 140$ eV \\citep{George1991}. The value of $R_{\\rm d}\\simeq 0.27$ found based on the observed reflection strength then implies $w_{\\rm d}\\sim 38$ eV, somewhat higher than the estimate based on the line fluxes, but still in an approximate agreement taking into account measurement errors and a number of assumptions we have made. On the other hand, $R_{\\rm t}\\simeq 0.24$ found above would explain only $w_{\\rm t}\\sim 30$ eV. However, we note that we have neglected the local absorber, which obviously also gives rise to an Fe K$\\alpha$ line component. Its characteristic $N_{\\rm H}$ of $\\sim 10^{23}$ cm$^{-2}$ (Table \\ref{comp_results}b) can readily explain the excess equivalent width of $\\sim 70$ eV (e.g., \\citealt{Makishima1986,Awaki1991}) of the constant line component with respect to that expected from the torus.  We note a number of caveats for our results. The standard accretion disc is flared, not flat \\citep{Shakura1973}. This, however, would have a relatively minor effect, changing somehow the distribution of the inclination angles. Our results indicate disc reflection with $R_{\\rm d}$ substantially less than unity, which implies the X$\\gamma$ source is not entirely above the disc. A likely geometry explaining it is a hot inner flow surrounded by a disc (e.g. \\citealt{Abramowicz1995,Narayan1995,Yuan2001}), in which case the incident radiation will have much larger incident angles (measured with respect to the axis of symmetry) than those assumed in the used model \\citep{Magdziarz1995}. The disc may be warped (e.g., \\citealt{Wijers1999}), which again would change the distribution of the incident angles. Furthermore, Compton scattering in the hot plasma above the disc reduces the observed reflection strength \\citep{Petrucci2001b}.  Then, we also used the slab geometry for the torus reflection. Thus, the obtained solid angle does not correspond to the actual angle subtended by the torus from the X-ray source. For example, \\citet{Murphy2009} considered a torus with a circular cross section subtending (as seen from the centre) a $2\\upi$ solid angle. They found that the reflection component in this case is several times weaker than that corresponding to a slab subtending the same angle. One obvious effect here is that, unlike the case of a slab, an observer sees only a fraction of the reflecting surface, without parts obscured by the torus itself. If we take this into account, our value of $R_{\\rm t}\\sim 0.2$ appears consistent with a torus subtending a $\\sim 2\\upi$ solid angle. We note, however, that even if the torus solid angle is formally $2\\upi$ or so, a large part of it will be shielded from the X-ray source by the accretion disc and the black hole itself, so the actual irradiated solid angle may be substantially lower. This effect was not taken into account in the geometrical model of \\citet{Murphy2009}. Furthermore, the cross section of the torus may be substantially different from circular, which may increase the observed reflection, as well as it may be clumpy \\citep*{Krolik1988,Nenkova2002}, which will decrease it. Still, our value of $w_{\\rm t}\\sim 100$ eV agrees with that from a torus with the column density of $N_{\\rm H}\\sim 10^{24}$ cm$^{-2}$ for $\\Gamma\\simeq 1.75$ \\citep{Murphy2009}, which provides an explanation for that equivalent width alternative, or additional, to that as being due to the line emission of the absorber (discussed above).  We have assumed that the torus reflection component is constant on the time scale of years. On the other hand, \\citet{Minezaki2004} found that the dusty torus inner boundary is at $\\simeq 0.04$ pc. Thus, a fraction of the torus reflection may vary on the corresponding time scale of $\\sim 50$ d. Furthermore, 0.04 pc corresponds to $\\simeq 2\\times 10^4 R_{\\rm g}$ (where $R_{\\rm g}\\equiv GM/c^2$ is the gravitational radius), where the accretion disc may still be present and join onto the torus. On the other hand, \\citet{Radomski2003} have constrained the torus outer boundary to $\\la 35$ pc, so the bulk of the torus reflection may still be constant over a time scale of years.  \\subsection{Absorber properties} \\label{absorber}  We compare our results with those based on \\sax, which also provided broad-band spectra, allowing to simultaneously determine absorber properties and the continuum. They were studied by \\citet{Puccetti2007} and \\citet{DeRosa2007}, who used the same absorber model as in this work. We find their results to be compatible with ours. The fully covering absorber has $N_{\\rm H} \\simeq (0.9$--$9.4)\\times 10^{22}$ cm$^{-2}$ for \\sax\\/ and $\\simeq (3.0$--$8.1)\\times 10^{22}$ cm$^{-2}$ in our case, see Table \\ref{comp_results}(b). For the partially covering absorber, $N_{\\rm H}^{\\rm p}\\simeq  (3.5$--$30.3)\\times 10^{22}$ cm$^{-2}$ (\\sax) and $\\simeq (9.5$--$30.2)\\times 10^{22}$ cm$^{-2}$ (this work). The covering fractions are also similar, $f^{\\rm p}\\simeq 0.34$--0.71 (\\sax) and 0.36--0.71 (this work). The only exception is the long-exposure \\sax\\/ observation of NGC 4151 in December 2001 showing a very low $N_{\\rm H}$ \\citep{DeRosa2007}.  We find an anticorrelation between the $N_{\\rm H}$ of both absorber components and the hard X-ray flux. We define the total column, $N_{\\rm H}=N_{\\rm H}^{i}+ f^{p} N_{\\rm H}^{p}$. For the bright state, we have $N_{\\rm H}/10^{22}$ cm$^{-2}$ = { 11.1$\\pm$3.5} (\\xmm\\/ X1+X2) and { 8.2$\\pm$4.9} (X3), for the medium state,  { 18.0$\\pm$2.8} (\\xmm\\/ X5), and for the dim state, { 23.1$\\pm$0.8} (\\xmm\\/ X4) and { 23.2$\\pm$0.9} (\\suzaku). Also, the covering fraction is anti-correlated with the hard X-ray flux, see Table \\ref{comp_results}(b). A similar trend is seen for the results of \\citet{Puccetti2007}, who used the 6--10 keV for the X-ray flux. The variable part of the absorber needs to be relatively close to the X-ray source to be able to follow its flux on the time scale of days. The absorber in NGC 4151 was identified with massive outflows from the accretion disc \\citep{Piro2005}, a broad-line region \\citep{Puccetti2007} or the surface or wind of the torus \\citep{Schurch2002}. Only the disc wind is close enough to the X-ray source to follow the change of the nuclear emission and to produce any correlation. The physical explanation for the correlation remains unclear; possibly the wind rate in NGC 4151 decreases with increasing accretion rate.  \\subsection{The nature of the X-ray source} \\label{xray}  The geometry and parameters of the source are constrained by a number of our findings. We find (i) that both the Compton parameter and the X-ray spectral index are approximately the same in both bright and and dim states (Sections \\ref{cutoff}--\\ref{Compton}). This implies an approximately constant amplification ratio of the Comptonization process, i.e., the ratio of the power emitted by the plasma to that supplied to it by seed photons (e.g., \\citealt{Beloborodov1999}). The amplification factors obtained with the Comptonization model are indeed similar, $A$ = { 17$\\pm$2, 15$\\pm$2} for the bright and dim state, respectively (assuming $kT_{\\rm bb}=10$ eV). A similar value of $A\\simeq 13$  was obtained for a Comptonization model fitted to the 1991 data from \\rosat, \\ginga\\/ and OSSE \\citep{Zdziarski1996a}.  If the seed photons are supplied by an accretion disc surrounding the hot plasma, this implies an approximately constant disc inner radius. Based on this, we have inferred (ii) a relatively weak reflection from the disc, $R_{\\rm d}\\simeq 0.3$ (Section \\ref{close_distant}). This qualitatively agrees with $A\\gg 1$ and both findings rule out the static disc corona geometry; the (outer) disc subtends a small solid angle as seen from the X-ray source, and the X-ray source subtends a small solid angle as seen from inner parts of the disc \\citep{Zdziarski1999}. An implication of the latter is that the modelled disc blackbody emission (which provides seed photons for Comptonization) should be much weaker than that observed (which corresponds to the entire disc emission). This seems to be indeed the case as shown in Fig.\\ \\ref{sed}, where we see that the bright-state UV flux inferred from the Comptonization model is about an order of magnitude below the shown maximum observed far UV (1350 \\AA; 9.2 eV) flux. On the other hand, the model dim-state far UV flux is close to the historical minimum observed. However, that minimum represented a single isolated dip in the light curve \\citep{Kraemer2006}, and the actual far UV flux corresponding to the dim state is likely to be significantly higher. In choosing the shown range of the far UV flux we used its strong correlation with the optical flux, which can  be seen by comparing Fig.\\ \\ref{longest}(d) with fig.\\ 1 in \\citet{Kraemer2006}.  The IR, optical and UV fluxes shown in Fig. \\ref{sed} appear relatively weakly affected by the host galaxy emission. In particular, the dominance of the AGN in the U band is shown by its strong variability \\citep{Czerny2003}. In the optical range, there can be some non-negligible fraction of emission from the broad and narrow line regions, but in the UV we expect that the disc dominates. Given that the measured UV emission is, furthermore, absorbed by the host galaxy, the conclusion above that only a small fraction of the disc emission undergoes Comptonization appears secure.  \\begin{figure} \\includegraphics[width=\\columnwidth]{ngc4151_sedis5eV.eps} \\caption{The broad-band spectrum of NGC 4151. The radio, IR, optical and UV data (black dots) were taken from the NED database. Additional data are shown in circles in radio (green, \\citealt{Ulvestad2005}), mid and near IR (brown, \\citealt{Radomski2003,Ruiz2003}, optical (magenta, OMC, this work) and UV (purple, \\swift/UVOT, this work; and cyan, \\hst/STIS, the flux extrema from fig.\\ 3 of \\citealt{Kraemer2006}). The \\xmm\\/ EPIC pn and ISGRI spectra of the bright and dim states are shown by the blue and red dots, respectively, the corresponding unabsorbed Comptonization models below 10 keV are shown by dashes, and the disc blackbody (assuming geometry parameter $=0$ and $kT_{\\rm bb}=5$ eV) incident on the hot plasma are shown by the dotted curves. The upper limits in the 0.75--1--3 MeV energy ranges are from COMPTEL \\citep{Maisack1995}.} \\label{sed} \\end{figure}  Another result is (iii) the connection of the normalization of the fitted disc blackbody seed emission to the disc inner radius, $R_{\\rm in}$. Based on the definition from {\\sc xspec}, we can express it as \\begin{equation} K={10^8 r\\cos i\\over f_{\\rm col}^4} {(R_{\\rm in}/10^{12}\\,{\\rm cm})^2\\over (D/10\\,{\\rm Mpc})^2}, \\label{norm} \\end{equation} where $f_{\\rm col}\\sim 1.7$ is the colour correction to the blackbody temperature \\citep{Shimura1995}, and $r$ is the ratio of the blackbody emission incident on the plasma to that emitted by the disc. It implies (at $D=13.2$ Mpc and $i=45\\degr$), \\begin{equation} R_{\\rm in}\\simeq 4.5\\times 10^{13} \\left(K \\over 10^{9}\\right)^{1/2} \\left( r\\over 0.1\\right)^{-1/2} \\left(f_{\\rm col}\\over 1.7\\right)^2\\, {\\rm cm}. \\end{equation} At $M=4.6\\times 10^7\\msun$ \\citep{Bentz2006}, $R_{\\rm g}\\simeq 6.8\\times 10^{12}$ cm. Thus, the values of $K$ in Table \\ref{comp_results}(a) for the bright state (at $T_{\\rm bb}=10$ eV) imply $R_{\\rm in}\\sim 6 R_{\\rm g}$, the last stable orbit for a non-rotating black hole. This is in conflict with the findings above, according to which the inner part of the flow is occupied by the hot plasma and not by the disc, unless the black hole is rotating fast. Furthermore, the {\\sc diskbb} model used for the seed photons becomes invalid when $R_{\\rm in}$ is close to the last stable orbit with the standard zero-stress boundary condition \\citep{Gierlinski1999}. We note, however, the disc flux is $\\propto R_{\\rm in}^2 T_{\\rm bb}^4$, and approximately $R_{\\rm in}\\propto T_{\\rm bb}^{-2}$ for a given X$\\gamma$ spectrum. In particular, $K=7.2\\times 10^9$ for the bright state at $kT_{\\rm bb}=5$ eV (consistent with the UV data) and the geometry parameter $=0$, confirming the above scaling. This yields $R_{\\rm in} \\simeq 18 R_{\\rm g}$, allowing the existence of a hot inner flow. Given that agreement, we used $kT_{\\rm bb}=5$ eV in the models shown in Fig.\\ \\ref{sed}. We note that if $R_{\\rm in}\\simeq$ constant, $kT_{\\rm bb}$ has to be somewhat lower in the dim state.  For comparison, the radius, $R_{\\rm eq}$, at which the integrated gravitational energy dissipation at $R>R_{\\rm eq}$ equals that at $R<R_{\\rm eq}$ is $R_{\\rm eq}\\simeq 14.7 R_{\\rm g}$ for a non-rotating black hole in the pseudo-Newtonian approximation \\citep{Paczynski1980} and with the zero-stress boundary condition at the last stable orbit, and less for a rotating black hole. Thus, $R_{\\rm in}\\sim R_{\\rm eq}$, as found above, is consistent with the observed energy distribution of NGC 4151, showing the optical/UV emission (presumably due to the disc) at a similar level as that at $\\sim 100$ keV (due to the hot flow). Note that $R_{\\rm in}$ moderately larger than $R_{\\rm eq}$ is required if some fraction of the energy dissipated below $R_{\\rm in}$ is advected to the black hole instead of being radiated away. However, a disc truncated at $R_{\\rm in}\\gg R_{\\rm eq}$ would be in strong conflict with the broad-band spectrum, as it implies the disc UV emission much weaker than that of X-rays. A prediction of this result is a presence of an Fe K$\\alpha$ line only moderately relativistically broadened (as found earlier by Z02).  The reflection of $R_{\\rm d} \\simeq 0.3$ may be easily achieved if the central hot region (with a large scale height) is surrounded by a disc. We have calculated the reflection strength in some simple geometrical models, with a sphere and cylinder surrounded by a flat disc. For the sphere and cylinder with the height equal to 1/2 of its radius we find $R_{\\rm d}\\simeq 0.29$ and 0.26, respectively, in agreement with our estimate. At this $R_{\\rm d}$, the emission of the parts of the disc close to the hot flow is most likely dominated by reprocessing of X-rays, in agreement with the constancy of the $y$ parameter \\citep{Zdziarski1999}. As a caveat, we point out that the geometrical $R_{\\rm d}$ estimates are highly simplified, not taking into account the actual hot flow geometry and the radial distribution of the Comptonized emission. We also note that the constancy of the $y$ parameter can also be explained in a model in which the inner disc radius increases in the dim state, but at the same time the energy dissipation in the hot flow becomes relatively stronger in its outer region, e.g., due to advection.  On the other hand, the effective solid angle subtended by the hot flow as seen from the disc can be identified with our parameter $r\\sim 0.1$ above. All the above considerations very strongly argue for the Comptonizing plasma in NGC 4151 being photon starved. This is not compatible with a static corona above the disc \\citep{Haardt1991}, and compatible with a hot inner flow surrounded by an outer cold disc. However, it is also possible that the corona does cover the disc but it is in a state of a mildly relativistic outflow \\citep*{Beloborodov1999,Malzac2001}. The relativistic motion then both reduces the flux of the up-scattered photons incident on the disc, which explains the relative weakness of the disc reflection, and reduces the flux of the disc blackbody photons entering the corona in its comoving frame, which explains the photon starvation of the hot plasma, i.e., $A\\gg 1$. A prediction of this model is the presence of a very broad relativistic Fe line (as the disc now extends all the way to the minimum stable orbit). The lack of detection of such a line in high-quality \\xmm\\/ and \\suzaku\\/ spectra represents an argument against this model.  Our next finding has been that (iv) the plasma temperature increased by a factor of $\\simeq 3.5$ when the bolometric X-ray luminosity decreased by about the same factor, see Table \\ref{comp_results}(a). In agreement with the approximate constancy of $y$, the optical depth also decreased by a similar factor. This can be explained if the accretion rate, $\\dot M$, is $\\propto L_{\\rm X\\gamma}$ and the hot flow has a velocity profile independent of $\\dot M$, so $\\tau\\propto L_{\\rm X\\gamma}$. The temperature then follows from the energy balance provided the ratio of the power in the hot flow emission to that in seed photons cooling remains constant. However, models of accreting hot flows yield a dependence of $\\tau$ on $L_{\\rm X\\gamma}$ weaker than the simple proportionality [e.g., eqs.\\ (10--11) in \\citealt{Z98}; table 2 in \\citealt{Yuan2004}].  The total model luminosity of the Comptonization component (including reflection) in the bright and dim state is $L_{\\rm X\\gamma}\\simeq 5.2 \\times 10^{43}$ erg s$^{-1}$ and $1.5\\times 10^{43}$ erg s$^{-1}$, respectively (see Table \\ref{comp_results}a). The estimated average luminosity from the IR to UV is $\\simeq 7 \\times 10^{43}$ erg s$^{-1}$, with the IR band ($<10^{14}$ Hz) contributing $\\sim 60$ per cent of it. If we then assume that the IR flux remains constant and the optical--UV flux scales as the high-energy emission, the total bolometric luminosity can be estimated as $1.3\\times 10^{44}$ erg s$^{-1}$, $7\\times 10^{43}$ erg s$^{-1}$ in the bright and dim state, respectively. Assuming the hydrogen mass fraction of $X=0.7$ and $M=4.6 \\times 10^{7} \\msun$, the Eddington luminosity, $L_{\\rm E}\\equiv 8\\upi GM m_{\\rm p} c/[(1+X)\\sigma_{\\rm T}]\\simeq 6.8 \\times 10^{45}$ erg s$^{-1}$, giving the Eddington ratios of $\\simeq 0.01$--0.02. Note the factor of 2 variability of the bolometric luminosity compared to the factor of 3.5 variability of the Comptonization component. This difference appears to be explained by the presence of the constant, averaged over very long time scales, IR component.  Relevant time scales include the light crossing time across the circumference of the last stable orbit, which for a non-rotating black hole is $\\sim 2700$ s (at the default mass). The corresponding light crossing time at $R_{\\rm in}\\sim 15 R_{\\rm g}$ is $\\sim 7000$ s. Thus, the hour-time scale variability observed by ISGRI in the bright state (Fig.\\ \\ref{short}) appears to correspond to fluctuations at those radii. On the other hand, the viscous time of an accretion flow is, \\begin{eqnarray} \\lefteqn{t_\\mathrm{visc}= {GM\\over c^3} \\left(R\\over R_{\\rm g}\\right)^{3/2} \\alpha^{-1} \\left(H\\over R\\right)^{-2} }\\\\ \\lefteqn{ \\qquad\\!\\! \\simeq 1.2\\times 10^6\\,{\\rm s} \\left(R\\over 15R_{\\rm g}\\right)^{3/2} \\left(M\\over 4.6\\times 10^7\\msun\\right) \\left(\\alpha\\over 0.1\\right)^{-1} \\left(H/R\\over 1/3\\right)^{-2},\\nonumber} \\label{eq:tvisc} \\end{eqnarray} where $H$ is the flow scale height, $\\alpha$ is the viscosity parameter, and $H/R\\sim 1/3$ for a hot flow. At our default mass, $15 R_{\\rm g}$ and $\\alpha=0.1$, $t_\\mathrm{visc}\\simeq 13$ d. For comparison, the \\swift/BAT data for NGC 4151 need to be rebinned to several days in order to get a sufficient signal-to-noise ratio for low flux states. Thus, this detector can sample the overall behaviour of the hot flow but not variability close to the last stable orbit.  The spectrum of NGC 4151 resembles spectra of black-hole binaries in the hard state, also commonly fitted by thermal Comptonization. However, a faint non-thermal tail has occasionally been seen there, e.g., \\citet{McConnell2002}, \\citet{Wardzinski2002}. The Comptonization model used by us also allows to include a corresponding non-thermal tail to the electron distribution. However, the data for NGC 4151 do not allow us to test for the presence of any non-thermal spectral component, see Section \\ref{Compton}.  \\subsection{Comparison with other objects} \\label{other}  We have presented the average parameters of Seyfert spectra in Section \\ref{intro}, for both e-folded power law and Comptonization models. Our results for NGC 4151 are within the range observed for other Seyferts, as noted by Z02. Z02 also provide a critical analysis of earlier finding of very hard X-ray spectra in NGC 4151, which they find to result from improper modelling of its complex absorption. NGC 4151 in the bright state is only slightly harder than the average Seyfert 1, and it shows less reflection, in accordance with the reflection-index correlation \\citep{Zdziarski1999}. On the other hand, the dim state appears to have a significant contribution to reflection from a distant reflector, possibly a torus (Section \\ref{close_distant}).  The similarity between spectra of Seyferts and black-hole binaries in the hard spectral state was pointed out by, e.g., \\citet{Zdziarski1996b}, though the latter appear to have on average somewhat harder spectra. This is likely to be due to the disc seed photon energies being much higher than those in Seyferts, which reduces the spectral index, $\\Gamma$, for a given Comptonization amplification factor, $A$ \\citep{Beloborodov1999}.  A striking similarity between the OSSE spectra of the black-hole binary GX 339--4 in the hard state and NGC 4151 was pointed out by \\citet{Zdziarski1998}. \\citet{Wardzinski2002} have shown that when $L_{\\rm X\\gamma}$ of GX 339--4 decreased by a factor of $\\sim 2$, $kT_{\\rm e}$ increased significantly (see their fig.\\ 1a), from $46^{+6}_{-4}$ keV to $76^{+5}_{-6}$ keV, with only a slight softening of the X-ray slope. This effect is very similar to that seen in NGC 4151, see Table \\ref{comp_results}(a). In both NGC 4151 and GX 339--4, this effect may be due to a decrease of $\\dot M$ in the hot flow causing the associated reduction of $\\tau$. ",
        "conclusions": "\\label{conclusions}  We have presented a comprehensive spectral analysis of all \\integral\\/ data obtained so far for NGC 4151, together with all contemporaneous data from \\xte, \\xmm, \\swift\\/ and \\suzaku\\/ (Section \\ref{data}). Our main findings are summarized below.  We have found that the 20--100 keV emission measured by \\integral\\/ has had almost the same range of fluxes as that measured by other satellites during past 40 years (Section \\ref{variability}). Thus, our analysis appears to explore the full range of the variability of this object. Also, we have found that this flux appears to uniquely determine the intrinsic broad-band spectrum.  Simultaneous observations by \\integral\\/ in the optical range and in medium and hard X-rays show a very strong correlation within the X-ray band, and a less clear correlation between the optical and X-ray emission, with a bimodal behaviour of the optical flux. The linearity of the former correlation shows that the X-ray spectral slope of NGC 4151 remains almost constant.  Most of the \\integral\\/ observations correspond to either a bright or dim hard X-ray state (Section \\ref{assumptions}). We have found that a thermal Comptonization model provides very good fits to the data (Section \\ref{Compton}). As the state changes from bright to dim (with $L_{\\rm X\\gamma}$ decreasing by a factor of $\\simeq 3.5$), the plasma electron temperature increases from $kT_{\\rm e}\\sim 60$ keV to $\\sim 200$ keV and the Thomson optical depth decreases from $\\tau\\sim 2$ to $\\sim 0.6$ (for spherical geometry), i.e., $\\tau\\propto 1/kT_{\\rm e}$ and $\\tau \\propto L_{\\rm X\\gamma}$. The former proportionality corresponds to an almost constant Compton $y$ parameter, with the X-ray slope remaining almost constant, with only a slight softening in the dim state, in agreement with the medium/hard X-ray correlation. This is suggestive of almost constant source geometry and amplification factor. The latter proportionality may occur due to the accretion rate varying at an approximately constant hot flow velocity.  The fitted strength of Compton reflection increased from $R\\simeq 0.4$ to $R\\simeq 0.8$, which, at the face value, would be in conflict with the constant source geometry indicated by the constant $y$. However, in accordance with the AGN unified model, NGC 4151 is likely to possess a remote torus, also reflecting the central X-ray emission. Assuming the solid angles subtended by both the (close) disc reflector and the (distant) torus reflector are constant, we can explain the varying fitted reflection strength by the presence of a contribution of the flux reflected by the torus, which is constant over the observation time scale given its very large size (Section \\ref{close_distant}). In the bright state, the contribution of the torus reflection is small, but it becomes much stronger in the dim state, explaining the fitted large value of $R$. Given the above assumption, we find that the solid angles subtended by the disc and torus (as seen from the X-ray source) are $\\simeq 0.3$, $0.2$ of $2\\upi$, respectively. These values have been obtained with a slab reflection model, which is likely to substantially underestimate the actual solid angle subtended by the torus, which might then be close to $2\\upi$.  The Comptonizing plasma is photon-starved, i.e., the flux in seed photons incident on the plasma is $\\sim$15 times weaker than the Comptonized flux (Section \\ref{xray}). This is consistent with the seed photons being supplied by a truncated optically-thick disc surrounding a hot accretion flow. This geometry is also consistent with the relatively small disc reflection. We also find that the flux of the disc photons fitted to the X-ray spectra is an order of magnitude below the UV fluxes actually observed. This is again consistent with the above geometry, in which only a small fraction of the emitted disc photons is incident on the inner hot regions. All these findings rule out a static disc corona geometry. However, they are still compatible with a corona outflowing at a mildly relativistic speed, see Section \\ref{xray}.  In any case, the disc inner radius cannot be very large given that the observed $EF_E$ at UV are of the same order as those at hard X-rays, implying a rough equipartition between the disc and hot flow emission. For a non-rotating black hole, the radius with the same integrated dissipation below and above it is $\\simeq 15 R_{\\rm g}$. The truncation radius can also be constrained by the fitted Comptonization model, which assumes the seed photons are disc blackbody. This gives a similar value of $R_{\\rm in}$ for the inner disc temperature of $\\sim 5$ eV, which value is consistent with the observations.  Compared to Seyfert 1s, NGC 4151 has a somewhat harder X-ray spectrum and less reflection. Its spectrum is also similar to black-hole binaries in the hard state, in particular to GX 339--4, which has shown a very similar spectral evolution.  Future missions, e.g., \\textit{Astro-H\\/} and \\textit{EXIST}, with instruments with a sensitivity comparable to that of OSSE, can provide better data above 200 keV than obtained up to date. Yet, for a better understanding of the physics of the central engine of AGNs, we need high-quality data up to at least 1 MeV. This would yield strong constraints on the electron distribution in the Comptonization region (e.g. the presence of a non-thermal tail) and on the source geometry (e.g. the spatial distribution of the electron temperature and the characteristic optical depth."
    },
    "1005/1005.3070_arXiv.txt": {
        "abstract": "{\\mia\\ is a Galactic \\hii\\ region with a bubble morphology, detected at optical and radio wavelengths. It is considered the first observational confirmation of the collect-and-collapse model of triggered star-formation. } {We aim to analyze the dust and gas properties of the \\mia\\ region to better constrain its effect on local future generations of stars. In addition, we investigate the relationship between the dust emissivity index $\\beta$ and the dust temperature, $T_{dust}$.} {Using \\textit{Herschel} PACS and SPIRE images at $100$, $160$, $250$, $350$ and $500\\,\\mu$m we determine $T_{dust}$ and $\\beta$ throughout \\mia, fitting the spectral energy distributions (SEDs) obtained from aperture photometry. With the SPIRE Fourier transform spectrometer (FTS) we obtained spectra at different positions in the \\mia\\ region. We detect $J$-ladders of $^{12}$CO and $^{13}$CO, with which we derive the gas temperature and column density. We also detect proxies of ionizing flux as the \\nii\\ and \\cia\\ transitions.} {We find an average value of $\\beta \\sim 1.5$ throughout \\mia, as well as a $T_{dust}$ difference between the photodissociation region (PDR, $\\sim25\\,$K) and the interior ($\\sim40\\,$K) of the bubble. We recover the anti-correlation between $\\beta$ and dust temperature reported numerous times in the literature. The relative isotopologue abundances of CO appear to be enhanced above the standard ISM values, but the obtained value is very preliminary and is still affected by large uncertainties. } {} ",
        "introduction": "Sharpless 104 (\\mia, \\citealt{sharpless1959}) is an optically visible Galactic \\hii\\ region with a bubble morphology, excited by an O6V star \\citep{crampton1978, lahulla1985}. It is located $\\sim4\\,$kpc from the Sun \\citep{deharveng2003}, with galactic coordinates $74.7620; +0.60$ (J2000).  \\citet{deharveng2003} proposed \\mia\\ as a strong candidate for massive triggered star formation through the collect-and-collapse process \\citep{elmegreen1977}. The ionized region is also visible at radio wavelengths \\citep{fich1993}, and an ultracompact (UC) \\hii\\ region, coincidant with the IRAS~20160+3636 source, lies at its eastern border (\\citealt{condon1998}).  We present new submm images and spectra taken towards \\mia\\ with the \\textit{Herschel} Space Observatory \\citep{pilbratt2010}. These observations allow us to map a wavelength range not easily accessed before, providing new insights into the dust and gas properties of \\mia. ",
        "conclusions": "With \\textit{Herschel} PACS and SPIRE data we have analysed dust and gas properties of the bubble-shaped \\hii\\ region \\mia. Aperture photometry of PACS+SPIRE images allowed us to derive the dust emissivity index $\\beta$ and the dust temperature throughout \\mia. We found two different groups, one at colder temperatures and higher $\\beta$, and the other at warmer temperatures and lower $\\beta$. We recover the inverse $\\beta - T_{dust}$ relationship reported in the literature, but the uncertainty in our fit prevents us from finding a precise depence. As an example, the different relationships found by \\citet{dupac2003} and \\citet{desert2008} are both compatible with our results. PACS and SPIRE have allowed us to constrain the ``cold'' side of the SED. We have estimated both $\\beta$ and $T_{dust}$ simultaneously, but the uncertainties remain relatively high. With the data at hand we can only fit a single temperature component, which disregards the contribution of warmer components to the PACS shorter wavelengths. The existance of a $\\beta - T_{dust}$ anti-correlation could be due to differences of the dust grain properties between the PDR and the \\hii\\ cavity.  SPIRE-FTS spectra at different pointings throughout \\mia\\ have unveiled the CO chemistry of the region in more detail. We detect the $^{12}$CO and $^{13}$CO $J$-ladders up to the $J = 13\\rightarrow12$ and $J = 9\\rightarrow8$ transitions respectively, revealing the warm gas component in the region. Rotational diagrams towards the UC\\hii\\ region in the PDR of \\mia\\ show that the $^{13}$CO emission is optically thin and also the $^{12}$CO for transitions above the $J=8$ level. The emission shows two different gas components, a colder one with a temperature of $\\sim170\\,$K and a hotter one at a temperature of $\\sim250\\,$K. The CO column densities derived would suggest an enhancement of the $^{13}$CO isotopologue abundance ratio with respect to the elemental value, but the uncertainties of the different assumptions are still too large to confirm that result. In a follow-up paper we will show models of the PDR of \\mia, which will provide better constraints on the gas temperature, density and column density structure."
    },
    "1005/1005.1844_arXiv.txt": {
        "abstract": "{ Active galactic nuclei (AGNs) have typically been discovered in massive galaxies of high metallicity.} {We attempt to increase the number of AGN candidates in low metallicity galaxies. We present VLT/UVES and archival VLT/FORS1 spectroscopic and NTT/SUSI2 photometric observations of the low-metallicity  emission-line galaxy Tol~2240--384 and perform a detailed study of its morphology, chemical composition, and emission-line profiles.} {The profiles of emission lines in the UVES and FORS1 spectra are decomposed into several components with different kinematical properties by performing multicomponent fitting with Gaussians. We determine abundances of nitrogen, oxygen, neon, sulfur, chlorine, argon, and iron by analyzing the fluxes of narrow components of the emission lines using empirical methods. We verify with a photoionisation model that the physics of the narrow-line component gas is similar to that in common metal-poor galaxies. } {Image deconvolution reveals two high-surface brightness regions in Tol~2240--384 separated by 2.4 kpc. The brightest southwestern region is surrounded by intense ionised gas emission that strongly affects the observed $B-R$ colour on a spatial scale of $\\sim$5 kpc. The profiles of the strong emission lines in the UVES spectrum are asymmetric and all these lines apart from H$\\alpha$ and H$\\beta$ can be fitted by two Gaussians of FWHM $\\sim$ 75 -- 92 km s$^{-1}$ separated by $\\sim$ 80 km s$^{-1}$ implying that there are two regions of ionised gas emitting narrow lines. The oxygen abundances in both regions are equal within the errors and in the range 12+log O/H = 7.83 -- 7.89. The shapes of the H$\\alpha$ and H$\\beta$ lines are more complex. In particular, the H$\\alpha$ emission line consists of two broad components of FWHM $\\sim$ 700 km s$^{-1}$ and 2300 km s$^{-1}$, in addition to narrow components of two regions revealed from profiles of other lines. This broad emission in H$\\alpha$ and H$\\beta$ associated with the high-excitation, brighter southwestern H {{\\sc ii}} region of the galaxy is also present in the archival low- and medium-resolution VLT/FORS1 spectra. The extraordinarily high luminosity of the broad H$\\alpha$ line of 3 $\\times$ 10$^{41}$ erg s$^{-1}$ cannot be accounted for by massive stars at different stages of their evolution. The broad H$\\alpha$ emission persists over a period of 7 years, which excludes supernovae as a powering mechanism of this emission. This emission most likely arises from an accretion disc around a black hole of mass $\\sim$ 10$^7$ $M_\\odot$.} {} ",
        "introduction": "Active galactic nuclei (AGNs) are understood to be powered by massive black holes at the centers of galaxies, accreting gas from their surroundings. They are usually found in massive, bulge-dominated galaxies and their gas metallicities are generally high \\citep{S98,H02,H09}. However, it remains unclear whether AGNs in low-metallicity low-mass galaxies do exist. \\citet{Gr06} and \\citet{B08} searched the Sloan Digital Sky Survey (SDSS) spectroscopic galaxy samples for low-mass Seyfert 2 galaxies. In particular, \\citet{Gr06} used a sample of 23 000 Seyfert 2 galaxies selected by \\citet{K03} and found only $\\sim$ 40 Seyfert 2 galaxies among them with masses lower than  10$^{10}$ $M_\\odot$. They demonstrated, however, that the metallicities of these AGNs are around solar or slightly subsolar. The same high metallicity range is found in the SDSS sample of 174 low-mass broad-line AGNs of \\citet{G07}. On the other hand, \\citet{I07} and \\citet{IT08} demonstrated that broad-line AGNs with much lower metallicities probably exist, although they occupy a region in the diagnostic diagram differing from that of more metal-rich AGNs and are extremely rare. They identified four of these galaxies, which were found to have oxygen abundances 12 + log O/H in the range 7.36 -- 7.99 on the basis of a systematic search for extremely metal-deficient emission-line dwarf galaxies in the SDSS Data Release 5 (DR5) database of 675 000 spectra. The absolute magnitudes of those four low-metallicity AGNs are typical of dwarf galaxies, their host galaxies have a compact structure, and their spectra resemble those of low-metallicity high-excitation H~{\\sc ii} regions. \\citet{I07} found that there is however a striking difference: the strong permitted emission lines, mainly the H$\\alpha$ $\\lambda$6563 line, show very prominent broad components characterised by properties unusual for dwarf galaxies: 1) their H$\\alpha$ full widths at zero intensity (FWZI) vary from 102 to 158 \\AA, corresponding to expansion velocities between 2200 and 3500 km s$^{-1}$; 2) the broad H$\\alpha$ luminosities $L_{br}$ are extraordinarily large, between  3$\\times$10$^{41}$ and 2$\\times$10$^{42}$ erg s$^{-1}$.  This is higher than the range 10$^{37}$--10$^{40}$ erg~s$^{-1}$ found by \\citet{I07} for the other emission-line galaxies (ELGs) with broad-line emission.  The ratio of H$\\alpha$ flux in the broad component to that in the narrow component varies from 0.4 to 3.4, compared to 0.01--0.20 for the other galaxies; 3) the Balmer lines exhibit a very steep decrement, which is indicative of collisional excitation and the broad emission originating in very dense gas ($N_e$$>$10$^{7}$ cm$^{-3}$).  To account for the broad-line emission in these four objects, \\citet{I07} considered various physical mechanisms such as Wolf-Rayet (WR) stars, stellar winds from Ofp or luminous blue variable stars, single or multiple supernova (SN) remnants propagating in the interstellar medium, and SN bubbles. While these mechanisms may be able to produce $L_{br}$ $\\sim$ 10$^{36}$ to 10$^{40}$ erg s$^{-1}$, they cannot generate yet higher luminosities, which are more likely associated with SN shocks or AGNs. \\citet{I07} considered type IIn SNe because their H$\\alpha$ luminosities are higher ($\\sim$10$^{38}$--10$^{41}$ erg~s$^{-1}$) than those of the other SN types and they decrease less rapidly. \\citet{IT08} found no significant temporal evolution of broad H$\\alpha$ in all four galaxies over a period of 3--7 years. Therefore, the IIn SNe mechanism may be excluded, leaving only the AGN mechanism capable of accounting for the high luminosity of the broad H$\\alpha$ emission. However, we also have difficulty with this mechanism. In particular, all four galaxies are present in neither the ROSAT catalogue of the X-ray sources nor the NVSS catalogue of radio sources. High-ionisation emission lines such as He {\\sc ii} $\\lambda$4686 or [Ne {\\sc v}] $\\lambda$3426 are weak or not detected in optical spectra. Based on the observational evidence, \\citet{IT08} concluded that all four studied galaxies most likely belong to the very rare type of low-metallicity AGNs in which non-thermal ionising radiation is strongly diluted by the radiation of a young massive stellar population.  The fifth galaxy of this type, Tol 2240--384, was first spectroscopically studied by \\citet{T91} and \\citet{M94}. However, in those low-resolution spectra, some important emission lines, such as [O~{\\sc ii}] $\\lambda$3727, [Ne~{\\sc iii}] $\\lambda$3868, and H$\\alpha$ $\\lambda$6563, are missing. This, in particular, precludes abundance determination and the detection of broad hydrogen emission. \\citet{K04,K06} studied Tol 2240--384 spectroscopically, and \\citet{K06} derived the oxygen abundance of this galaxy, 12+logO/H = 7.77 $\\pm$ 0.08. We note that no broad emission was reported by \\citet{T91}, \\citet{M94}, and \\citet{K06}. In this paper, we present 8.2m Very Large Telescope (VLT) spectroscopic observations and 3.5 ESO New Technology Telescope (NTT) photometric observations of this emission-line galaxy. Its optical spectrum shows the very broad components of hydrogen emission lines and is similar to those found previously by \\citet{I07} and \\citet{IT08} for the four other galaxies. We describe observations in Sect.2. The morphology of the galaxy is discussed in Sect.3 and its location in the emission-line diagnostic diagram is discussed in Sect.4. Element abundances are derived in Sect.5. The kinematics of the ionised gas from narrow emission lines is discussed in Sect.6. We discuss in Sect.7 the properties of the broad emission and derive the mass of the central black hole assuming an AGN mechanism for the origin of the broad line emission. Our conclusions are summarized in Sect.8.     \\begin{figure*}[t] \\hspace*{0.5cm}\\psfig{figure=14390f1.ps,angle=-90,width=16.cm,clip=} \\caption{Flux-calibrated VLT/UVES spectrum of Tol 2240--384, obtained on 23 August 2009, corrected for the redshift of $z$ = 0.07595 [ESO program 383.B-0271(A)] (upper spectrum). The lower spectrum is the upper spectrum downscaled by a factor of 100. The scale of the ordinate is that for the upper spectrum. Note the broad emission in the hydrogen line H$\\alpha$ $\\lambda$6563. No appreciable broad emission is detected in other strong permitted and forbidden lines, which is indicative of the rapid motions of relatively dense ionised gas with an electron number density $N_e$ $\\geq$ 10$^7$ cm$^{-3}$.} \\label{fig1} \\end{figure*}  \\begin{figure*}[t] \\hspace*{1.0cm}\\psfig{figure=14390f2a.ps,angle=-90,width=7.5cm}% \\hspace*{1.0cm}\\psfig{figure=14390f2b.ps,angle=-90,width=7.5cm,clip=} \\caption{Flux-calibrated and redshift-corrected archival VLT/FORS1 medium-resolution (left) and low-resolution (right) spectra of Tol 2240--384 obtained on 12 September 2002 [ESO program 69.C-0203(A)] (upper spectra). The lower spectra are the upper spectra downscaled by a factor of 100. The scale of the ordinate is that for the upper spectra. Note the strong broad emission in the hydrogen line H$\\alpha$ $\\lambda$6563 and much weaker broad emission in the hydrogen line H$\\beta$ $\\lambda$4861. No appreciable broad emission is detected in strong forbidden lines.} \\label{fig2} \\end{figure*} ",
        "conclusions": "  1. Image deconvolution reveals two high-surface brightness regions in Tol~2240--384 separated by 2.4 kpc and differing in their luminosity by a factor of $\\sim$10. The brightest southwestern region is surrounded by intense ionised gas emission, which strongly affects the observed $B-R$ colour on a spatial scale of $\\sim$5 kpc. This high-excitation H~{\\sc ii} region is associated with broad H$\\alpha$ and H$\\beta$ emission. Surface photometry does not indicate, in agreement with the results of image deconvolution and unsharp masking, the presence of a bulge in Tol~2240--384.   2. We derived the oxygen abundance 12+logO/H = 7.85$\\pm$0.01 in Tol 2240--384, which is consistent within the errors with the value of 7.77$\\pm$0.08 derived earlier by \\citet{K06}.  3. The emission line profiles in the high resolution UVES spectrum reveal the presence of two narrow components with a radial velocity difference of $\\sim$ 78 km s$^{-1}$. Furthermore, the full widths at half maximum (FWHMs) of the narrow lines differ. Strong forbidden nebular lines [O~{\\sc iii}] $\\lambda$4959, 5007, [O~{\\sc ii}] $\\lambda$3726, 3729, and [Ne~{\\sc iii}] $\\lambda$3868 have FWHMs of 73 -- 80 km s$^{-1}$. The FWHMs of hydrogen lines are larger, $\\sim$ 85 km s$^{-1}$ and decrease from H$\\alpha$ to H$\\delta$ emission lines. The largest FWHM of $\\sim$ 92 km s$^{-1}$ is found for the He {\\sc i} $\\lambda$5876 emission line. This data suggest that narrow permitted hydrogen and helium lines probe the denser inner parts of the emitting regions compared to the forbidden lines.  4. Both UVES and FORS1 spectra reveal the presence of very broad hydrogen lines with FWHMs greater than 2000 km s$^{-1}$. The steep Balmer decrement of the broad hydrogen lines and the very high luminosity of the broad H$\\alpha$ line 3$\\times$10$^{41}$ erg s$^{-1}$ suggest that the broad emission arises from very dense and high luminosity regions such as those associated with supernovae of type IIn or with accretion discs around black holes. However, the presence of the broad H$\\alpha$ emission over a period of 7 years rules out the SN mechanism. Thus, the emission of broad hydrogen lines in Tol 2240--384 is most likely associated with an accretion disc around a black hole.  5. There is no obvious spectroscopic evidence of a source of non-thermal hard ionising radiation in Tol 2240--384. However, none is expected if, as we argue, the density of the broad line region is 5$\\times$10$^{6}$ -- 5$\\times$10$^{8}$ cm$^{-3}$.  6. Assuming that the broad emission in Tol 2240--384 is powered by an AGN, we have estimated a mass for the central black hole of $M_{\\rm BH}$ $\\sim$ 10$^7$ $M_\\odot$."
    },
    "1005/1005.4699_arXiv.txt": {
        "abstract": "We present an analysis of an 8 arcminute diameter map of the area around the galaxy cluster Abell 1835 from jiggle map observations at a wavelength of 1.1 mm using the Bolometric Camera (Bolocam) mounted on the Caltech Submillimeter Observatory (CSO). The data is well described by a model including an extended Sunyaev-Zel'dovich (SZ) signal from the cluster gas plus emission from two bright background submm galaxies magnified by the gravitational lensing of the cluster. The best-fit values for the central Compton value for the cluster and the fluxes of the two main point sources in the field: SMM J140104+0252, and SMM J14009+0252 are found to be $y_{0}=(4.34\\pm0.52\\pm0.69)\\times10^{-4}$, 6.5$\\pm{2.0}\\pm0.7$ mJy and 11.3$\\pm{1.9}\\pm1.1$ mJy, where the first error represents the statistical measurement error and the second error represents the estimated systematic error in the result. This measurement assumes the presence of dust emission from the cluster's central cD galaxy of $1.8\\pm0.5$~mJy, based on higher frequency observations of Abell 1835. The cluster image represents one of the highest-significance SZ detections of a cluster in the positive region of the thermal SZ spectrum to date. The inferred central intensity is compared to other SZ measurements of Abell 1835 and this collection of results is used to obtain values for $y_{0} = (3.60\\pm0.24)\\times10^{-4}$ and the cluster peculiar velocity $v_{z} = -226\\pm275$~km/s. ",
        "introduction": "The Sunyaev Zel'dovich (SZ) effect \\citep{Sunyaev70} is the redistribution of energy in the Cosmic Microwave Background (CMB) spectrum due to interations between CMB photons and hot electrons along the line of sight between the surface of last scattering and an observer. The main source for the SZ effect is from the hot gas that exists in the intra-cluster medium (ICM) of massive galaxy clusters \\citep{Birkinshaw99, Carlstrom02, Rephaeli06}. The SZ actually comprises two effects. The thermal effect consists of a dimming or decrement in the apparent brightness of the CMB towards a galaxy cluster at low frequencies and a corresponding brightening or increment at high frequencies with the null crossover point at approximately 215 GHz. The kinematic effect has the spectral dependence of a standard temperature shift in the CMB which has the same sign at all frequencies. The two effects can therefore be distinguished from each other with measurements at multiple frequencies.  Measurements of the amplitude of the SZ thermal distortion towards a cluster can be combined with measurements of the X-ray emission to determine the angular diameter distance, $d_{A}$, whose value depends on cosmology \\citep{Birkinshaw91}. Estimates of the Hubble constant have been made using this technique for a number of clusters (e.g. \\cite{Jones95, Grainge96, Holzapfel97, Tsuboi98, Mauskopf00, Reese03, Battistelli03, Udomprasert04}) and can be used to constrain cosmological models. In addition, because the SZ surface brightness for a cluster with a given mass is almost independent of redshift, SZ surveys can give information about the evolution of the number counts of clusters vs. redshift. This depends strongly on the evolution of the so-called dark energy or cosmological constant and therefore SZ surveys have been identified as one of the key probes of the nature of dark energy (e.g. \\cite{Diego02, Weller02, DeDeo05, Albrecht06, Bhattacharya07}). A number of dedicated SZ surveys are already producing results, in particular the Atacama Cosmology Telescope (ACT) and the South Polar Telescope (SPT) \\citep{Plagge09, Hincks09, High10}.  Most of the SZ detections reported to date have been made at low frequencies corresponding to the SZ decrement. Follow-up photometric or spectroscopic measurements of known clusters at higher frequencies would significantly improve the precision in the measurement of the kinematic SZ effect as well as constraining possible contamination in the low frequency data. Measurements at high frequencies corresponding to the SZ increment suffer from confusion from emission from dusty galaxies, including background high redshift galaxies amplified by the gravitational lensing of the cluster as well as increased atmospheric contamination from ground-based telescopes. Accurate measurement of the SZ increment requires a combination of angular resolution sufficient to resolve and remove point sources combined with high sensitivity and control of systematics necessary to detect the more diffuse SZ signal.  This paper presents analysis of observations of the galaxy cluster Abell 1835 using the Bolometric Camera (Bolocam) mounted on the Caltech Submillimeter Observatory (CSO), situated on the summit of Mauna Kea, Hawaii. We also compare the results of this analysis with other sets of data taken at different wavelengths. Abell 1835 is one of the most luminous clusters observed in the ROSAT catalogue and is well-known as a cooling core cluster. It is also known to contain two lensed sub-mm point sources, SMM J14009+0252 and SMM J140104+0252 (see e.g. \\citet{Ivinson00, Zemcov07}).  The paper is organized as follows: Section 2 describes the observations with the Bolocam instrument; Section 3 describes the analysis pipeline developed for processing the data; Section 4 describes the modelling of the data and the determination of the characteristic parameters of the model (as well as the errors in their values), and Section 5 discusses the results of the analysis of the Bolocam data and their combination with other literature results. ",
        "conclusions": "The observation and analysis of 1.1 mm Bolocam jiggle-map data of Abell 1835, including details of the process of converting the raw data to maps of the cluster, have been discussed. A parameter search has been used to determine the best fit values of the central Compton parameter for the cluster and the fluxes for the point sources SMM J140104+0252 and SMM J14009+0252. These values are found to be $y_{0}=(4.68\\pm0.48\\pm0.82)\\times 10^{-4}$, 6.5$\\pm{2.0}$ mJy and 11.3$\\pm{1.9}$ mJy, respectively with a statistical signal-to-noise of 10.5, 3.3 and 6.0 respectively. If the model is refined further to include a point source, representing dust emission from the cD galaxy in Abell 1835, the Compton parameter for the cluster is found to be $y_{0}=(4.34\\pm0.52\\pm0.69)\\times 10^{-4}$ The value for $y_{0}$ was compared to other literature results and found to agree well. The SZ spectrum for Abell 1835, based upon these values for $y_{0}$ was fit with the full SZ form, including both thermal and kinetic effects and relativistic corrections up to 7th order. The values of $y_{0}$ and $v_z$ which optimized this fit were found to be $(3.60\\pm0.24)\\times 10^{-4}$ and $-226\\pm 275$~km/s, which are both in agreement with previous literature results. It is also concluded that, in order to better evaluate the contamination from dust and accurately characterize the spectrum of point sources in the same field as SZ clusters more results need to be obtained at shorter wavelengths in order to obtain more a detailed spectral energy distribution."
    },
    "1005/1005.3587_arXiv.txt": {
        "abstract": "Because of an old quasar APM $08279+5255$ at $z=3.91$, some dark energy models face the challenge of the cosmic age problem. It has been shown by Wei and Zhang [Phys. Rev. D {\\bf 76}, 063003 (2007)] that the holographic dark energy model is also troubled with such a cosmic age problem. In order to accommodate this old quasar and solve the age problem, we propose in this Letter to consider the interacting holographic dark energy in a non-flat universe. We show that the cosmic age problem can be eliminated when the interaction and spatial curvature are both involved in the holographic dark energy model. ",
        "introduction": "\\label{sec1} The fact that our universe is undergoing accelerated expansion has been confirmed by lots of astronomical observations such as type Ia supernovae (SNIa)~\\cite{Riess98}, large scale structure (LSS)~\\cite{Tegmark04} and cosmic microwave background (CMB) anisotropy~\\cite{Spergel03}. It is the most accepted idea that this cosmic acceleration is caused by some kind of negative-pressure matter known as dark energy whose energy density has been dominative in the universe. The combined analysis of cosmological observations indicates that the universe today consists of about 70\\% dark energy, 30\\% dust matter (cold dark matter plus baryons), and negligible radiation. The famous cosmological constant $\\lambda$ introduced first by Einstein is the simplest candidate for dark energy. However, the cosmological constant scenario has to face the so-called ``fine-tuning problem'' and ``cosmic coincidence problem''~\\cite{dereview}. Many dark energy models have been proposed, while the nature of dark energy is still obscure. Besides quintessence~\\cite{quintessence}, a wide variety of scalar-field dark energy models have been studied including $k$-essence~\\cite{kessence}, hessence~\\cite{hessence}, phantom~\\cite{phantom}, tachyon~\\cite{tachyon}, quintom~\\cite{quintom}, ghost condensate~\\cite{ghost}, etc. In addition, there are other proposals on dark energy such as interacting dark energy models~\\cite{intde}, brane world models~\\cite{brane}, Chaplygin gas models~\\cite{cg}, Yang-Mills condensate models~\\cite{YMC}, and so on.  The dark energy problem is essentially an issue of quantum gravity, owing to the concern of the vacuum expectation value of some quantum fields in a universe governed by gravity. However, by far, we have no a complete theory of quantum gravity yet. So, it seems that we have to consider the effects of gravity in some effective quantum field theory in which some fundamental principles of quantum gravity could be taken into account. It is commonly believed that the holographic principle \\cite{hp} is just a fundamental principle of quantum gravity. Based on the effective quantum field theory, Cohen et al. \\cite{r16} pointed out that the quantum zero-point energy of a system with size $L$ should not exceed the mass of a black hole with the same size, i.e., $L^3\\Lambda^4\\leq LM_{Pl}^2$, where $\\Lambda$ is the ultraviolet (UV) cutoff of the effective quantum field theory, which is closely related to the quantum zero-point energy density, and $M_{Pl}\\equiv 1/\\sqrt{8\\pi G}$ is the reduced Planck mass. This observation relates the UV cutoff of a system to its infrared (IR) cutoff. When we take the whole universe into account, the vacuum energy related to this holographic principle can be viewed as dark energy (its energy density is denoted as $\\rho_{\\Lambda}$ hereafter). The largest IR cutoff $L$ is chosen by saturating the inequality, so that we get the holographic dark energy density \\begin{equation} \\label{eq1}\\rho_{\\Lambda}=3c^2M_{Pl}^2L^{-2} \\end{equation} where $c$ is a numerical constant characterizing all of the uncertainties of the theory, and its value can only be determined by observations. If we take $L$ as the size of the current universe, say, the Hubble radius $H^{-1}$, then the dark energy density will be close to the observational result. However, Hsu \\cite{r17} pointed out that this yields a wrong equation of state for dark energy. Subsequently, Li \\cite{Li04} suggested to choose the future event horizon of the universe as the IR cutoff of this theory. This choice not only gives a reasonable value for dark energy density, but also leads to an accelerated universe. Moreover, the cosmic coincidence problem can also be explained successfully in this model, provided that the inflation lasts for more than 60 $e$-folds. Most recently, a calculation of the Casimir energy of the photon field in a de Sitter space is performed \\cite{Li:2009pm}, and it is a surprising result that the Casimir energy is indeed proportional to the size of the horizon (the usual Casimir energy in a cavity is inversely proportional to the size of the cavity), in agreement with the holographic dark energy model.    Up to now, the holographic dark energy model has been tested by various observational data including SNIa~\\cite{r38}, SNIa+BAO+CMB~\\cite{holodata,Li:2009zs}, X-ray gas mass fraction of galaxy clusters~\\cite{r40}, differential ages of passively evolving galaxies~\\cite{r41}, Sandage-Leob test~\\cite{r42}, and so on \\cite{Shen:2004ck}. These analyses show that the holographic dark energy model is consistent with the observational data. However, Wei and Zhang~\\cite{hao} used some old high redshift objects (OHROs) to test the holographic dark energy model and found that the original holographic dark energy model can be ruled out unless a lower Hubble constant (e.g., $h=0.56$) is taken. So, according to Ref.~\\cite{hao}, there is a cosmic age crisis in the holographic dark energy model.  In fact, many dark energy models are in the face of such a cosmic age problem. In history, the cosmic age problem has been focused in cosmology for several times. At present, the cosmic age crisis coming from some OHROs appears again in cosmological models, even though dark energy is involved in the models. In cosmology there is a very basic principle that the universe cannot be younger than its constituents. So, if the age of some astronomical object (at some redshift) is measured accurately, then it can be used to test cosmological models according to this simple age principle. Now, there are some OHROs discovered, for example, the $3.5$ Gyr old galaxy LBDS 53W091 at redshift $z=1.55$~\\cite{r46} and the $4.0$ Gyr old galaxy LBDS 53W069 at redshift $z=1.43$~\\cite{r48}. In particular, the old quasar APM $08279+5255$ at redshift $z=3.91$ is an important one, which has been used as a ``cosmic clock'' to constrain cosmological models. Its age is estimated to be $2.0-3.0$ Gyr \\cite{r51}. These three OHROs at $z=1.43$, 1.55 and 3.91 have been used to test many dark energy models, including the $\\Lambda$CDM model~\\cite{r44}, the general EoS dark energy model~\\cite{r57}, the scalar-tensor quintessence model~\\cite{r58}, the $f(R)=\\sqrt{R^2-R_0^2}$ model~\\cite{r59}, the DGP braneworld model~\\cite{r60}, the power-law parameterized quintessence model~\\cite{r62}, the Yang-Mills condensate model~\\cite{Tong:2009mu}, the holographic dark energy model~\\cite{hao}, the agegraphic dark energy model~\\cite{Zhang:2007ps}, and so on. These investigations show that the two OHROs at $z=1.43$ and 1.55 can be easily accommodated in most dark energy models, whereas the OHRO at $z=3.91$ cannot, even in the $\\Lambda$CDM model~\\cite{r44} and the holographic dark energy model ~\\cite{hao}.  In this Letter, we revisit the cosmic age problem in the holographic dark energy model. We consider an interacting holographic dark energy model in a non-flat universe. We will show that the age crisis in the original holographic dark energy model can be avoided when the interaction and the spatial curvature are involved in the holographic dark energy model. ",
        "conclusions": "\\label{sec4}  In this Letter, we have revisited the cosmic age problem in the holographic dark energy model. The cosmic age problem brought by the old quasar APM $08279+5255$ has caused trouble to many cosmological models, and the holographic dark energy model is not an exception either \\cite{hao}. In order to accommodate the old quasar APM $08279+5255$ in the holographic dark energy model, we propose to consider the interaction between dark energy and matter in the model. We have shown that the quasar indeed can be accommodated in the holographic dark energy model when an appropriate interaction strength is chosen. Taking the current observational constraints \\cite{Li:2009zs} into account, we have demonstrated that both interaction and spatial curvature should be simultaneously involved in the holographic dark energy model. It has been shown that if such a sophisticated case is considered the quasar APM $08279+5255$ can be accommodated and the cosmic age problem can thus be avoided in the holographic dark energy model. The price of solving the age problem in this way is also apparent, i.e., the model involves too many free parameters, which may weaken the plausibility of the model, to some extent.  It is well known that the consideration of interaction in the holographic dark energy can be used to avoid the future big-rip singularity caused by $c<1$ \\cite{Li:2009zs,Li:2008zq}. In this Letter we have provided another advantage for the consideration of interaction in the holographic dark energy, i.e., the interaction between dark energy and matter can also be used to avoid the age problem caused by the old quasar. So, our result can be viewed as a further support to the interacting holographic dark energy model.  Of course, we have to confess that the age problem would still exist if some extreme cases are taken into account, say, a much larger possible age of the quasar $t_{obj}$ with a larger $h$. It is remarkable that the age of the old quasar APM $08279+5255$ has not been measured accurately yet, and the age problem caused by this quasar has troubled many dark energy models (including the $\\Lambda$CDM model). It is expected that the future accurate measurement on the age of this old quasar would eliminate the cosmic age crisis in dark energy models."
    },
    "1005/1005.5642_arXiv.txt": {
        "abstract": "In a pilot project to study the relationship between star formation and molecular gas properties in nearby normal early-type galaxies, we have obtained observations of dense molecular gas tracers in the four galaxies of the {\\tt SAURON} sample with the strongest $^{12}$CO emission.  We used the Institut de Radio Astronomie Millimetrique (IRAM) 30m telescope $3$ and $1$~mm heterodyne receivers to observe $^{13}$CO(J=$1$--$0$), $^{13}$CO(J=$2$--$1$), HCN(J=$1$--$0$) and HCO$^+$(J=$1$--$0$). We report the detection of $^{13}$CO emission in all four {\\tt SAURON} sources and HCN emission in three sources, while no HCO$^+$ emission was found to our detection limits in any of the four galaxies.  We find that the $^{13}$CO/$^{12}$CO ratios of three {\\tt SAURON} galaxies are somewhat higher than those in galaxies of different Hubble types. The HCN/$^{12}$CO and HCN/$^{13}$CO ratios of all four {\\tt SAURON} galaxies resemble those of nearby Seyfert and dwarf galaxies with normal star formation rates, rather than those of starburst galaxies. The HCN/HCO$^+$ ratio is found to be relatively high (i.e., $>$1) in the three {\\tt SAURON} galaxies with detected HCN emission, mimicking the behaviour in other star-forming galaxies but being higher than in starburst galaxies. When compared to most galaxies, it thus appears that $^{13}$CO is enhanced (relative to $^{12}$CO) in three out of four {\\tt SAURON} galaxies and HCO$^+$ is weak (relative to HCN) in three out of three galaxies.  All three galaxies detected in HCN follow the standard HCN--infrared luminosity and dense gas fraction--star formation efficiency correlations. As already suggested by $^{12}$CO observations, when traced by infrared radiation, star formation in the three {\\tt SAURON} galaxies thus appears to follow the same physical laws as in galaxies of different Hubble types. The star formation rate and fraction of dense molecular gas however do not reach the high values found in nearby starburst galaxies, but rather resemble those of nearby normal star-forming galaxies. ",
        "introduction": "\\label{sec:intro}  \\begin{figure*} \\begin{center} \\resizebox{\\hsize}{!}{\\rotatebox{-90}{\\includegraphics{sauron-spec4.eps}}} \\caption{Emission line spectra of the four {\\tt SAURON} early-type galaxies observed. From top to bottom in each panel: HCO$^+$(J=1--0), HCN(J=1--0), $^{13}$CO(J=1--0) and $^{13}$CO(J=2--1).  From left to right: NGC~3032, NGC~4150, NGC~4459 and NGC~4526.  The $^{13}$CO emission is overlaid with the spectra of the $^{12}$CO emission ({\\it blue}) from \\citet{com07}, scaled down to match the $^{13}$CO line emission; the scaling factors are indicated in the respective spectra. The (single, dual, triple) Gaussian fits to the line spectra are plotted in red.} \\label{fig:fig1} \\end{center} \\end{figure*}  Over the last twenty years, growing evidence has accumulated that nearby early-type galaxies (E/S0s) are not always devoid of molecular gas, and can even harbour a substantial amount of it \\citep*[e.g.,][]{lee91,wik95,bet03,sag07,com07,you09}, often settled in a regularly rotating disc \\citep*[e.g.,][]{you02,you05,you08,cro08,cro09,cro10}. Interestingly, not all early-type galaxies with a substantial amount of molecular gas show obvious signs of (current) star formation \\citep[SF; e.g.,][]{cro08,cro09,cro10,you08}.  However, the converse is generally true \\citep[e.g.,][]{jeo09}, and the $\\ga30$~per cent fraction of SF objects within the nearby early-type galaxy population \\citep[e.g.,][]{yi05} appears to be superficially consistent with their $\\approx25$~per cent CO detection rate \\citep[e.g.,][]{sag06,sag07,com07,you09}. This fact raises many as yet unanswered questions: how is the (current) SF ignited in early-type galaxies? Is SF coupled to the molecular discs, i.e.\\ do all molecular gas discs form stars? Does SF only take place in the densest regions, perhaps ill-traced by the $^{12}$CO lines used in most surveys?  How, if at all, does the morphology of the molecular discs and the physical characteristics of the molecular gas correlate with the ages of the stars? In order to answer some of these questions, it is essential to study the spatially- and kinematically-resolved properties of the stellar populations, as well as those of the molecular gas in early-type galaxies (E/S0).  CO emission can also easily trace signs of gravitational interactions or other dynamical perturbations that can help to understand how star formation is induced.  Observations of a sample of 48 nearby E/S0 galaxies have been carried out with a custom-designed panoramic optical integral-field spectrograph \\citep{bac01} within the framework of the {\\tt SAURON} project. A description of the sample and illustrative results can be found in \\citet{zee02}. Results include the detection of kinematic misalignments, twists and decoupled cores as well as central discs in roughly half of the sources \\citep{ems04,kra08}. Information on the ionized gas (distribution, kinematics and ionisation; \\citealt{sar06,sar10}) and the age, metallicity and alpha-element enhancement of the stellar populations \\citep[][]{kun06,kun10} are also available. The specific angular momentum of the galaxies appears to be a dominant factor in their evolution \\citep{ems07,cap07}, and all CO detections are found in fast-rotating systems \\citep{you09}.  In order to study star formation in the {\\tt SAURON} sources and understand the various phenomena observed in the stellar and ionised gas data, observations looking for central molecular gas in the nearby {\\tt SAURON} E/S0 galaxies were obtained by \\citet{com07}.  $^{12}$CO emission was detected in $12$ of $43$ galaxies observed with the Institut de Radio Astronomie Millimetrique (IRAM) 30m telescope, suggesting typical gas masses in a range between a few 10$^{6}$ to 10$^{8}$ M$_\\odot$.  \\citet{com07} further found that the amount of molecular gas in the galaxies correlates with their far-infrared (FIR) luminosities, an often-used indicator of star formation. The gas-rich sources show the most pronounced star formation, as for normal star-forming spiral galaxies. However, the lack of correlations between the molecular gas and most stellar properties suggests that the molecular gas is largely unrelated to the old, pre-existing stellar population, and might have been externally accreted, as a result for instance of a merger or interaction event \\citep{com07}.    \\begin{table} \\caption{Basic parameters of the four {\\tt SAURON} galaxies.} \\label{tab:tab1} \\centering \\begin{tabular}{lrrrrc} \\hline NGC  & RA(J2000)      & DEC(J2000) & $V_\\odot$    & D       & Activity \\\\ & (hh:mm:ss.s) & (dd:mm:ss) & (km~s$^{-1}$) & (Mpc) & Type     \\\\ \\hline 3032 & 09:52:08.2 & $+$29:14:10 & 1533 & 21.7 & H~{\\small II}         \\\\ 4150 & 12:10:33.6 & $+$30:24:06 &  226 & 13.7 & none        \\\\ 4459 & 12:29:00.0 & $+$13:58:43 & 1210 & 16.3 & H~{\\small II}+L   \\\\ 4526 & 12:34:03.0 & $+$07:41:57 &  550 & 16.3 & none        \\\\ \\hline \\end{tabular}  Notes: Coordinates, distances, velocities and activity types are taken from the NASA/IPAC Extragalactic Database (NED). Galaxy types: L=LINER, H~{\\small II}=star formation bursts. \\end{table}  \\begin{table*} \\caption{Properties of the molecular line spectra.} \\label{tab:tab2} \\centering \\begin{tabular}{lrrrrrr} \\hline \\\\[-0.2cm] Line   & $^{12}$CO(J=1--0)  & $^{12}$CO(J=2--1) & $^{13}$CO(J=1--0)  & $^{13}$CO(J=2--1) & HCN(J=1--0) & HCO$^+$(J=1--0)$^b$ \\\\ Parameters$^a$   & (CYB07) & (CYB07) & (this paper) & (this paper) & (this paper) & (this paper)\\\\[0.2cm] \\hline \\\\[-0.2cm] NGC3032: & \\\\ \\hskip 0.4cm $I$ (K~km~s$^{-1}$)         & 9.3$\\pm$0.3 & 7.9$\\pm$0.2  & 0.9$\\pm$0.1 & 1.4$\\pm$0.2 & 0.32$\\pm$0.09 & $<$0.2\\\\ \\hskip 0.4cm $<v_\\odot>$ (km~s$^{-1}$)   & 1549$\\pm$2  & 1537$\\pm$2   & 1571$\\pm$10 & 1561$\\pm$10 & 1571$\\pm$13   &  \\\\ \\hskip 0.4cm $\\Delta v$ (km~s$^{-1}$)    & 129$\\pm$4   & 103$\\pm$3    &  130$\\pm$14 &  104$\\pm$10 &  160$\\pm$30   &  \\\\ \\\\[-0.2cm] NGC4150: & \\\\ \\hskip 0.4cm $I$ (K~km~s$^{-1}$)         & 6.1$\\pm$0.5 & 13.2$\\pm$0.5 & 0.4$\\pm$0.1 & 1.1$\\pm$0.2 & $<$0.2  & $<$0.2 \\\\ \\hskip 0.4cm $<v_\\odot>$ (km~s$^{-1}$)   & 204$\\pm$7   & 200$\\pm$3    & 221$\\pm$14  & 226$\\pm$12  &               &  \\\\ \\hskip 0.4cm $\\Delta v$ (km~s$^{-1}$)    & 158$\\pm$14  & 189$\\pm$8    & 150$\\pm$30  & 160$\\pm$30  &               &  \\\\ \\\\[-0.2cm] NGC4459: & \\\\ \\hskip 0.4cm $I$ (K~km~s$^{-1}$)         & 10.9$\\pm$0.5& 14.3$\\pm$0.5 & 3.2$\\pm$0.2   & 4.9$\\pm$0.4 & 0.8$\\pm$0.1 & $<$0.7\\\\ \\hskip 0.4cm $<v_\\odot>$ (km~s$^{-1}$)   & 1169$\\pm$9  & 1176$\\pm$8   & 1194$\\pm$22   & 1179$\\pm$40 & 1140$\\pm$20 &  \\\\ \\hskip 0.4cm $\\Delta v$ (km~s$^{-1}$)    & 477$\\pm$16  & 405$\\pm$15   & 390$\\pm$20    & 380$\\pm$20  & 440$\\pm$70  &  \\\\ \\\\[-0.2cm] NGC4526: & \\\\ \\hskip 0.4cm $I$ (K~km~s$^{-1}$)         & 23.8$\\pm$0.8& 37.4$\\pm$0.7 & 6.0$\\pm$0.2   & 11$\\pm$0.4 & 1.6$\\pm$0.1 & $<$0.7\\\\ \\hskip 0.4cm $<v_\\odot>$ (km~s$^{-1}$)   & 697$\\pm$2   & 695$\\pm$5    & 605$\\pm$30    & 600$\\pm$30  & 606$\\pm$40  & \\\\ \\hskip 0.4cm $\\Delta v$ (km~s$^{-1}$)    & 650$\\pm$23  & 533$\\pm$11   & 687$\\pm$26    & 677$\\pm$27  & 760$\\pm$50  & \\\\ \\\\[-0.2cm] \\hline \\end{tabular}  Notes: $^a$Line properties have been obtained by fitting a single Gaussian to the data from \\citet{com07} and this paper; $I\\equiv\\int~T_{\\rm mb}~dv$ is the velocity-integrated line intensity; $<v_\\odot>$ is the mean (central) heliocentric velocity; $\\Delta v$ is the velocity width (FWHM). $^b$ Determined from 3$\\sigma$ upper limits from spectrum at $\\sim$100~km~s$^{-1}$ spectral resolution assuming the same FWHM as found for the HCN(J=1--0) or $^{13}$CO(J=1--0) line. \\end{table*}  However, as $^{12}$CO emission has been found to be a rather unreliable tracer of the dense gas where star formation takes place \\citep[e.g.,][]{gao04a,gao04b}, as a pilot study we focus here on the volume and/or column density tracers HCN, HCO$^+$ and $^{13}$CO in the four {\\tt SAURON} galaxies with the brightest $^{12}$CO emission. Basic parameters of these four sources are given in Table~\\ref{tab:tab1}.  They form an interesting subset of the {\\tt SAURON} sample as they nicely populate the two main evolutionary groups: i) those with a more 'troubled' recent past, likely involving accretion and/or minor interactions with other galaxies (NGC~3032 and NGC~4150); and ii) those with a more peaceful immediate past dominated by secular evolution (NGC~4459 and NGC~4526). The first group appears to harbor a large fraction of widely distributed young stars while the second group possesses a lower young stellar fraction that is spatially limited to the center \\citep[][]{kun10}. Additionally, NGC~3032 shows signs that its young stars have a lower metallicity than its bulk stellar population. Nevertheless, only NGC~3032 and NGC~4459 show signs of circumnuclear activity both in the form of H~{\\small II} regions, indicative of star formation bursts, and, in the case of NGC~4459 also in the form of a low-ionization nuclear emission region (LINER) nucleus.  The paper is organised as follows. Section~\\ref{sec:obs} presents the observations. Section~\\ref{sec:results} consists of the results and a discussion of the data and their trends. Finally, section~\\ref{sec:conclusions} summarizes our findings and conclusions.  \\begin{table*} \\caption{Molecular line ratios for a representative set of nearby galaxies covering the Hubble sequence and different activity types.} \\label{tab:tab3} \\begin{center} \\begin{tabular}{lcccccccc} \\hline Source & Activity Type$^a$ &  $\\frac{^{13}{\\rm CO(J=1-0)}}{^{12}{\\rm CO(J=1-0)}}$ & $\\frac{^{13}{\\rm CO(J=2-1)}}{^{12}{\\rm CO(J=2-1)}}$ & $\\frac{{\\rm HCN(J=1-0)}}{^{13}{\\rm CO(J=1-0)}}$ & $\\frac{{\\rm HCN(J=1-0)}}{^{12}{\\rm CO(J=1-0)}}$ & $\\frac{{\\rm HCN(J=1-0)}}{{\\rm HCO^+(J=1-0)}}$ & Ref$^b$  & Region$^c$ \\\\%[0.15cm] \\hline \\multicolumn{9}{c}{Ellipticals: $-$6 $\\leq$ T  $\\leq$ $-$4}\\\\[0.1cm] Cen~A    &   Sy2 & $\\sim$0.07 &  $\\sim$0.08  & $\\sim$0.9 & $\\sim$0.06 & $\\sim$0.6$^d$ & (1) & Centre \\\\ & & $\\sim$0.1  &  $\\sim$0.09  & $\\sim$0.2 & $\\sim$0.02 &  --           & (1) & Dust Lane \\\\[0.2cm] \\multicolumn{9}{c}{Sauron (S0): $-$4 $<$ T $\\leq$ 0}\\\\[0.1cm]  NGC~3032 &   H~{\\small II}   & 0.09$\\pm$0.01 & 0.20$\\pm$0.02 & 0.36$\\pm$0.07 & 0.03$\\pm$0.01 & $>$1.7 & (2) & Centre \\\\ NGC~4150 &   NONE  &  0.07$\\pm$0.01 & 0.08$\\pm$0.01 &  $<$0.5   &  $<$0.03          & --     & (2) & Centre \\\\ NGC~4459 &   H~{\\small II}+L &  0.30$\\pm$0.02 & 0.34$\\pm$0.02 & 0.23$\\pm$0.03 & 0.07$\\pm$0.01 & $>$1.1 & (2) & Centre \\\\ NGC~4526 &   NONE  &  0.26$\\pm$0.01 & $\\sim$0.3        & 0.35$\\pm$0.02 & 0.09$\\pm$0.01 & $>$2.3 & (2) & Centre \\\\[0.15cm] \\multicolumn{9}{c}{Seyferts: 0 $<$ T $<$ +8}\\\\[0.1cm]  NGC~1068$^\\star$  &   Sy2        & 0.09$\\pm$0.02   & 0.06$\\pm$0.01   & 1.6$\\pm$0.6    & 0.13$\\pm$0.02   & 1.7$\\pm$0.2  & (3,6) & Centre\\\\ &            & $\\sim$0.02      & 0.062$\\pm$0.002 & $\\sim$80       &  $\\sim$1.4      & $\\sim$1.4    & (4,19) & CND$^d$\\\\ NGC~2237          &   ?          & $\\sim$0.1       & --              & 0.41$\\pm$0.07  & 0.04$\\pm$0.01   & --          & (4,5)   & Centre \\\\ NGC~3079          &   Sy2/L      &   0.06$\\pm$0.01   &   0.07$\\pm$0.01   &   0.6$\\pm$0.2    &   0.039$\\pm$0.09  &   $>$5        &   (6,23)  &   Centre \\\\ NGC~3627          &   Sy2/L      & $\\sim$0.06      & --              & $\\sim$1.1      & $\\sim$0.07      & 1.0$\\pm$0.1 & (3,8)   & Centre \\\\ NGC~3628          &   L+H~{\\small II}      &   0.12$\\pm$0.02   &   0.10$\\pm$0.02   &   0.34$\\pm$0.09  &   0.027$\\pm$0.007 &   1.6$\\pm$0.5 &   (22,26) &   Centre \\\\ NGC~3982          &   Sy1.9+H~{\\small II}  & 0.07$\\pm$0.01   & --              & 0.32$\\pm$0.08  & 0.022$\\pm$0.005 & --          & (4)     & Centre \\\\ NGC~4051          &   Sy1/NL     & 0.06$\\pm$0.01   & --              & 0.32$\\pm$0.08  & 0.019$\\pm$0.004 & --          & (4)     & Centre \\\\ NGC~4258          &   Sy1.9/L    & 0.10$\\pm$0.01   & --              & 0.29$\\pm$0.03  & 0.029$\\pm$0.003 & --          & (4)     & Centre \\\\ NGC~4388          &   Sy2        & 0.03$\\pm$0.01   & --              & 0.3$\\pm$0.1    & 0.010$\\pm$0.003 & --          & (4)     & Centre \\\\ NGC~4826          &   Sy2+H~{\\small II}    & 0.12$\\pm$0.02   & 0.13$\\pm$0.02   & 0.40$\\pm$0.04  & 0.06$\\pm$0.01   & 1.7$\\pm$0.1 & (3,6)   & Centre\\\\ NGC~4945          &   Sy2        &   0.06$\\pm$0.01   &   0.11$\\pm$0.01   &   0.75$\\pm$0.02  &   0.05$\\pm$0.01   &    $\\sim$1 &   (27,28) &   Centre \\\\ NGC~5033$^\\star$  &   Sy1.8      & 0.12$\\pm$0.01   & --              & 0.36$\\pm$0.03  & 0.044$\\pm$0.005 & $\\sim$1.9   & (4)     & CND$^d$ \\\\ NGC~5194$^\\star$  &   Sy2        & $\\sim$0.1-0.2   & $\\sim$0.1       & $\\sim$2        & $\\sim$0.5       & 1.4$\\pm$0.1 & (3,7,8) & CND$^d$ \\\\ NGC~6951$^\\star$  &   Sy2/L      &  --             & 0.13$\\pm$0.02   & $\\gtrsim$2     & $\\gtrsim$2      & 1.4$\\pm$0.1 & (3,5)   & CND$^d$ \\\\ NGC~7172          &   Sy2+H~{\\small II}    & $\\sim$0.1-0.2   & --              & 0.33$\\pm$0.07  & $\\sim$0.03      & --          & (4)     & Centre \\\\ NGC~7314          &   Sy1.9      & $\\sim$0.1-0.2   & --              & $<$0.4         & $\\sim$0.05      & --          & (4)     & Centre \\\\ NGC~7331          &   L          &   $\\sim$0.14      &   $\\sim$0.17      &   $\\sim$0.03     &   $\\sim$0.04      &   $>$1.7        &   (22,24) &   Center \\\\ NGC~7582          &   Sy2        &   --              &   --              &   --             &   0.03$\\pm$0.01   &   1.3$\\pm$0.6   &   (29)    &   Centre \\\\ NGC~7469$^\\star$     &   Sy1.2   &   $\\sim$0.06  &   $\\sim$0.05      &   0.95$\\pm$0.1   &   $\\sim$0.1       &   $\\sim$1.4     &   (4,6,7) &   Centre \\\\ Mrk~231$^{\\dag\\dag}$ &   Sy1     &   --          &   $<$0.03         &   --             &   $\\sim$0.2       &   0.72$\\pm$0.09 &   (29,30) &   Centre \\\\ Mrk~331$^\\dag$       &   Sy2+H~{\\small II} &   --          &   --              &   --             &   $\\sim$0.2       &   $\\sim$1.3     &   (29)    &   Centre \\\\ IRAS~05414+5840$^\\dag$ &   Sy2   &   --          &   --              &   --             &   0.040$\\pm$0.006 &   0.6$\\pm$0.2   &   (29)    &   Centre\\\\[0.15cm]   \\multicolumn{9}{c}{Starbursts: 0 $<$ T $<$ +8}\\\\[0.1cm] M82               &   SB+H~{\\small II}       & $\\sim$0.03-0.1 & $\\sim$0.1           & $\\sim$1       & $\\sim$0.2      & 0.7$\\pm$0.1   & (3,8,9)   & Centre \\\\ M83               &   SB+H~{\\small II}       &   0.10$\\pm$0.01  &   $\\sim$0.10$\\pm$0.03 &   $\\sim$1       &   $\\sim$0.1      &   1.3$\\pm$0.1   &   (23,25)   &   Centre \\\\ NGC~253           &   SB+Sy2+H~{\\small II}   & $\\sim$0.1      & $\\sim$0.1-0.3       & $\\sim$0.2-1.0 & $\\sim$0.05-0.3 & 0.8$\\pm$0.1   & (10,11)   & Centre \\\\ IRAS~210293       &   SB?          &   0.06$\\pm$0.01  &   0.11$\\pm$0.02       &      --\t  &      --\t       &   --            &   (20)      &   Centre \\\\ NGC~660           &   SB?+Sy2/L+H~{\\small II} &   0.07$\\pm$0.01  &   0.05$\\pm$0.01      &   0.7$\\pm$0.3   &   0.05$\\pm$0.02  &   1.0$\\pm$0.3   &   (20,21)   &   Centre \\\\ NGC~891\t          &   H~{\\small II}          &   $\\sim$0.2      &   --                  &   --            &   0.020$\\pm$0.005&   1.1$\\pm$0.5   &   (29,31)   &   Centre \\\\ NGC~986           &   SB+H~{\\small II}       &   0.10$\\pm$0.01  &   0.06$\\pm$0.01       &   0.9$\\pm$0.2   &   0.09$\\pm$0.02  &   -             &   (20)      &   Centre \\\\ NGC~1808          &   Sy2          &   0.06$\\pm$0.01  &   0.07$\\pm$0.01       &   1.0$\\pm$0.2   &   0.06$\\pm$0.01  &   0.56$\\pm$0.04 &   (20,21)   &   Centre \\\\ NGC~2146          &   SB+H~{\\small II}       &   0.08$\\pm$0.01  &   0.10$\\pm$0.01       &   0.8$\\pm$0.2   &   0.06$\\pm$0.01  &   0.77$\\pm$0.05 &   (4,20,21) &   Centre \\\\ NGC~2369$^\\dag$   &   SB?          &   --             &   --                  &   --            &   0.04$\\pm$0.01  &   1.0$\\pm$0.4   &   (29)      &   Centre \\\\ NGC~2903\t  &   SB?+H~{\\small II}      &   --             &   --                  &   --            &   0.020$\\pm$0.005&   0.4$\\pm$0.1   &   (29)      &   Centre \\\\ NGC~3256          &   SB+H~{\\small II}       &   0.03$\\pm$0.01  &   0.10$\\pm$0.04       &   2.0$\\pm$0.4   &   0.06$\\pm$0.01  &   --            &   (20)      &   Centre \\\\ NGC~4355\t  &   Sy2          &   --             &   --                  &   --            &   0.8$\\pm$0.2    &   0.08$\\pm$0.01 &   (29)      &   Centre  \\\\ NGC~6946          &   SB+Sy2+H~{\\small II}   & $\\sim$0.05-0.1 & --                  & $\\sim$1       & $\\sim$0.2      & 1.2$\\pm$0.1   & (3,4,8)   & Centre \\\\ NGC~7552          &   SB+L+H~{\\small II}     &   0.07$\\pm$0.01  &   0.11$\\pm$0.02       &   1.1$\\pm$0.2   &   0.08$\\pm$0.01  &   0.9$\\pm$0.2   &   (20,21)   &   Centre \\\\ \\hline \\end{tabular} \\end{center} Notes: T is the numerical Hubble type. Sources marked with a ``$\\star$'' show a particular gas chemistry with unusually high HCN/$^{12}$CO, HCN/$^{13}$CO and/or HCN/HCO$^+$ ratios; sources marked with a ``\\dag'' or ``\\dag\\dag'' are LIRGs and ULIRGs respectively. Values marked with a ``$\\sim$'' or given as a range indicate either values averaged over (or a range of values for) several positions/observations or values for which an accurate estimate of the uncertainty was not possible. $^a$ Sy=Seyfert; L=LINER, SB= starburst, H~{\\small II}=star formation bursts/H~{\\small II} regions, classifications taken from NED; $^b$ References: (1) \\citet*{wil00}, \\citet{wild97} and \\citet{wik97}; (2) this paper; (3) \\citet{kri08}; (4) Krips et al.\\ (2010, in prep.); (5) \\citet{pet03}; (6) \\citet{isr09a}; (7) \\citet*{isr06}; (8) \\citet{pag01}; (9) \\citet{mat00}; (10) \\citet{sak06}; (11) \\citet{sor01}; (12) \\citet{wils97}; (13) \\citet{tos07}; (14) \\citet{bro05}; (15) \\citet{pet98}; (16) \\citet{bol05}; (17) \\citet{hei99}; (18) \\citet{chin97}; (19) \\citet{gar08}; (20) \\citet{aalto95}; (21) \\citet{baan08} (and references therein); (22) \\citet{mei08}; (23) \\citet{ngrieu92}; (24) \\citet{isr99}; (25) \\citet{isr01}; (26) \\citet{isr09b}; (27) \\citet{hen94}; (28) \\citet{wan04}; (29) \\citet{baan08}; (30) \\citet{glen01}; (31) \\citet{saka97}; (32) \\citet{wil08}. $^c$ Regions in which the line ratios where determined: Centre = central $\\la10$~kpc of a galaxy; CND = circumnuclear disc, i.e.\\ a radius $<1$~kpc; GMC = average over several giant molecular clouds; SA = spiral arms. $^d$ Ratio determined from absorption lines. $^e$Ratios determined from interferometric observations. \\end{table*}  \\begin{table*} \\contcaption{} \\begin{center} \\begin{tabular}{lcccccccc} \\hline Source & Activity Type$^a$ & $\\frac{^{13}{\\rm CO(J=1-0)}}{^{12}{\\rm CO(J=1-0)}}$ & $\\frac{^{13}{\\rm CO(J=2-1)}}{^{12}{\\rm CO(J=2-1)}}$ & $\\frac{{\\rm HCN(J=1-0)}}{^{13}{\\rm CO(J=1-0)}}$ & $\\frac{{\\rm HCN(J=1-0)}}{^{12}{\\rm CO(J=1-0)}}$ & $\\frac{{\\rm HCN(J=1-0)}}{{\\rm HCO^+(J=1-0)}}$ & Ref$^a$  & Region$^b$ \\\\%[0.15cm] \\hline  \\multicolumn{9}{c}{Starbursts: 0 $<$ T $<$ +8}\\\\[0.1cm] NGC~7771$^\\dag$ &   SB+H~{\\small II}       &   --             &   --                  &   --            &   $\\sim$0.05     &   $\\sim$1.0     &   (29)      &   Centre \\\\ UGC~2855        &   ?            &   0.08$\\pm$0.02  &   0.11$\\pm$0.01       &       --        &       --         &   --            &   (20)      &   Centre \\\\ IC~860$^\\dag$          &   H~{\\small II}          &   --      &   --         &   -- &   0.08$\\pm$0.01 &   0.5$\\pm$0.2   &   (29) &   Centre \\\\ Arp~220$^{\\dag\\dag}$   &   SB+Sy2/L+H~{\\small II} &   $<$0.05 &   $\\sim$0.05 &   -- &   $\\sim$0.08    &   0.46$\\pm$0.09 &   (29) &   Centre \\\\ IRAS~22025+4205$^\\dag$ &   SB?          &   --      &   --         &   -- &   0.15$\\pm$0.03 &   0.42$\\pm$ 0.2 &   (29) &   Centre \\\\[0.15cm]  \\multicolumn{9}{c}{SF-spirals: 0 $<$ T $<$ +8}   \\\\[0.1cm] M33       &   H~{\\small II} & $\\sim$0.10    & $\\sim$0.14 &  --           & --              & --          & (12)    & GMC in SA\\\\ M31       &   L?  & $\\sim$0.13    & --         & $\\sim$0.15    & $\\sim$0.02      & $\\sim$0.9   & (13,14) & GMC in SA\\\\ Maffei~2  &   NONE  &   0.02-0.1      &   --         &   $\\sim$0.4-1.3 &   $\\sim$0.02-0.13 &   2.4$\\pm$0.4 &   (22,23) &   GMC in SA\\\\[0.15cm]  \\multicolumn{9}{c}{Dwarf galaxies: +8 $\\leq$ T $\\leq$ +10 }\\\\[0.1cm] IC10 &   SB   & --         & $\\sim$0.1     &  --           & --              & --            & (15)       & GMC \\\\ LMC  &   NONE & $\\sim$0.1  & $\\sim$0.2     & $\\sim$0.2-0.4 & $\\sim$0.03-0.06 & $\\sim$0.3-0.7 & (16,17,18) & GMC \\\\ SMC  &   NONE & $\\sim$0.1  & $\\sim$0.1-0.2 & $\\sim$0.3-0.4 & $\\sim$0.02      & $\\sim$0.4     & (16,17,18) & GMC \\\\ [0.15cm]  \\multicolumn{9}{c}{Peculiar galaxies and/or mergers}\\\\[0.1cm] NGC~6240$^\\dag$      &   Sy2/L  &   $\\sim$0.02 &   $\\sim$0.01 &   --     &   $\\sim$0.09      &   $\\sim$0.8   &   (29)    &   Centre \\\\ Mrk~273$^{\\dag\\dag}$ &   Sy2/L  &   --         &   --         &   --     &   $\\sim$0.2       &   $\\sim$ 0.4  &   (29)    &   Centre \\\\[0.2cm]  NGC~3620             &   SB?    &   --         &   --         &   --     &   0.08$\\pm$0.03   &   $<$0.4      &   (29)    &   Centre \\\\ IC~1623$^\\dag$       &   ?      &   --         &   --         &   --     &   0.040$\\pm$0.006 &   1.2$\\pm$0.2 &   (29)    &   Centre \\\\ Arp~55$^\\dag$        &   L+H~{\\small II}  &   --         & $<$0.04        &   --     &   $\\sim$0.1       &   $\\sim$0.8   &   (29,32) &   Centre \\\\ Arp~193$^\\dag$       &   L+H~{\\small II}  &   --         &   --         &   --     &   $\\sim$0.03      &   $\\sim$1.9   &   (29)    &   Centre \\\\ Arp~299A$^\\dag$      &   H~{\\small II}    &   $\\sim$0.04 &   $\\sim$0.03 &  $\\sim$1 &   0.040$\\pm$0.006 &   1.6$\\pm$0.3 &   (29)    &   Centre \\\\ Arp~299B$^\\dag$      &   H~{\\small II}    &   $\\sim$0.04 &   $\\sim$0.03 &  $\\sim$1 &   0.050$\\pm$0.004 &   $\\sim$0.8   &   (29)    &   Centre \\\\ IRAS~12112+0305$^\\dag$&   L+H~{\\small II} &   --         &   --         &   --     &   $\\sim$0.07      &   $\\sim$0.6   &   (29)    &   Centre \\\\ IRAS~15107+0724      &   H~{\\small II}    &   --         &   --         &   --     &   0.11$\\pm$0.03   &   0.5$\\pm$0.2 &   (29)    &   Centre \\\\  \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "\\label{sec:conclusions} We detected significant $^{13}$CO in four out of four and HCN emission in three out of four {\\tt SAURON} early-type galaxies, while no HCO$^+$ emission was found in any of the four sources. We find some pronounced differences in the line ratios of the {\\tt SAURON} galaxies when compared to other nearby galaxies of different Hubble and activity types. In particular, $^{13}$CO/$^{12}$CO appears slightly enhanced in three {\\tt SAURON} galaxies compared to other galaxy types. This may indicate different molecular gas excitation conditions and/or chemistry in these sources.  The closest resemblance to our four nearby {\\tt SAURON} galaxies is found in nearby star-forming galaxies, including Seyferts and dwarfs. This and the pronounced differences with starburst galaxies suggest that the four {\\tt SAURON} galaxies exhibit star-formation rates and efficiencies more similar to quiescent, normal star-forming galaxies than to starburst galaxies, and that they do not reach the high dense molecular gas fraction found in starbursts. Also, according to the high ($>$1) HCN/HCO$^+$ ratios, PDRs and/or supernovae explosions do not seem to play an important role in the chemistry of the molecular gas in our four targets.  The three {\\tt SAURON} galaxies observed in HCN nicely follow the same physical laws concerning star formation and dense molecular gas as other galaxies, at least when infrared radiation is used as the star formation tracer.  Although the four {\\tt SAURON} galaxies studied in this paper form a nice subset of the {\\tt SAURON} sample with respect to their general characteristics, they are not equally representative with respect to their molecular gas properties.  They have been chosen for their strong $^{12}$CO emission and could thus in principal be extreme cases within the sample. Follow-up observations of the $^{13}$CO, HCN and HCO$^+$ emission in a much larger sample of nearby early-type galaxies are currently underway based on the success of this pilot study. These will eventually allow us to confirm (or discard) the trends found in this pilot project and put them on a sounder statistical basis."
    },
    "1005/1005.3702_arXiv.txt": {
        "abstract": "{Observations of very high energy $\\gamma$-rays from blazars provide information about acceleration mechanisms occurring in their innermost regions. Studies of variability in these objects allow a better understanding of the mechanisms at play.  } {To investigate the spectral and temporal variability of VHE ($>100\\,\\mathrm{GeV}$) $\\gamma$-rays of the well-known high-frequency-peaked BL\\,Lac object \\pks\\ with the H.E.S.S. imaging atmospheric Cherenkov telescopes over a wide range of flux states.} {Data collected from 2005 to 2007 are analyzed. Spectra are derived on time scales ranging from 3 years to 4 minutes. Light curve variability is studied through doubling timescales and structure functions, and is compared  with red noise process simulations.} {The source is found to be in a low state from 2005 to 2007, except for a set of exceptional flares which occurred in July 2006. The quiescent state of the source is characterized by an associated mean flux level of $(4.32\\pm0.09_\\mathrm{stat}\\pm0.86_\\mathrm{syst}) \\times 10^{-11}\\,{\\rm cm^{-2}}\\,{\\rm s^{-1}}$ above $200\\,{\\rm GeV}$, or approximately $15\\%$ of the Crab Nebula, and a power law photon index of $\\Gamma=3.53\\pm0.06_\\mathrm{stat}\\pm0.10_\\mathrm{syst}$. During the flares of July 2006, doubling timescales of $\\sim 2\\,{\\rm min}$ are found. The spectral index variation is examined over two orders of magnitude in flux, yielding different behaviour at low and high fluxes, which is a new phenomenon in VHE $\\gamma$-ray emitting blazars. The variability amplitude characterized by the fractional r.m.s. $F_{\\rm var}$  is strongly energy-dependent and is $\\propto E^{0.19\\pm0.01}$. The light curve r.m.s. correlates with the flux. This is the signature of a multiplicative process which can be accounted for as a red noise with a Fourier index of $\\sim 2$.} {This unique data set shows evidence for a low level $\\gamma$-ray emission state from \\pks, which possibly has a different origin than the outbursts. The discovery of the light curve lognormal behaviour might be an indicator of the origin of aperiodic variability in blazars.}   \\offprints{santiago.pita@apc.univ-paris7.fr and francesca.volpe@mpi-hd.mpg.de/volpe@llr.in2p3.fr} ",
        "introduction": "\\label{intro}  The BL Lacertae (BL Lac) category of Active Galactic Nuclei (AGN) represents the vast majority of the population of energetic and extremely variable extragalactic very high energy $\\gamma$-ray emitters. Their luminosity varies in unpredictable, highly irregular ways, by orders of magnitude, and at all wavelengths across the electromagnetic spectrum. The very high energy (VHE, $E\\geq100\\,{\\rm GeV}$) $\\gamma$-ray fluxes vary often on the shortest timescales that can be seen in this type of object, with large amplitudes which can dominate the overall output. It hence indicates that the understanding of this energy domain is the most important one for understanding the underlying fundamental variability and emission mechanisms at play in high flux states.  It has been, however, difficult to ascertain whether $\\gamma$-ray emission is present only during high flux states or  also when the source is in a more stable or quiescent state but with a flux which is  below the instrumental limits. The advent of the current generation of atmospheric Cherenkov telescopes with unprecedented sensitivity in the VHE regime gives new insights into these questions.  The high frequency peaked BL Lac object (HBL) \\pks, located at a redshift $z=0.117$, initially discovered as a VHE $\\gamma$-ray emitter by the Mark 6 telescope (\\cite{cha99}), has been detected by the first H.E.S.S. telescope in 2002-2003 (\\cite{HESS2155_2003}~2005b).  It has been frequently observed by the full array of four telescopes since 2004, either sparsely during the H.E.S.S. monitoring program, or intensely during dedicated campaigns such as that described in \\cite{HESS2155_MWL}~(2005c), showing mean flux levels of $\\sim20\\%$ of the Crab Nebula flux for energies above $200\\,{\\rm GeV}$. During the summer of 2006, \\pks\\ exhibited unprecedented flux levels accompanied by strong variability (\\cite{HESS2155_BigFlare}~2007a), making temporal and spectral variability studies possible on timescales of the order of a few minutes. The VHE $\\gamma$-ray emission is usually thought to originate from a relativistic jet, emanating from the vicinity of a Supermassive Black Hole (SMBH). The physical processes at play are still poorly understood, but the analysis of the $\\gamma$-ray flux spectral and temporal characteristics is well suited to provide better insights.  For this goal the data set of H.E.S.S. observations of \\pks\\ between 2005 and 2007 is used. After describing the observations and the analysis chain in Section~\\ref{ObsAna}, the emission from the ``quiescent'', i.e. nonflaring, state of the source will be characterized in Section~\\ref{QuiescentState}. Section~\\ref{SpecVar} details spectral variability related to the source intensity. Section~\\ref{temporal_variab} will focus on the description of the temporal variability during the highly active state of the source, and its possible energy dependence. Section~\\ref{log_normal_process} will illustrate a description of the observed variability phenomenon by a random stationary process, characterized by a simple power density spectrum. Section~\\ref{characteristic_time} will show how limits on the characteristic time of the source can be derived. The multi-wavelength aspects from the high flux state will be presented in a second paper. ",
        "conclusions": "This data set, which exhibits unique features and results, is the outcome of a long-term monitoring program and dedicated, dense, observations. One of the main results here is the evidence for a VHE $\\gamma$-ray quiescent-state emission, where the variations in the flux are found to have a lognormal distribution. The existence of such a state was postulated by \\cite{ste96} in order to explain the extragalactic $\\gamma$-ray background at 0.03--$100\\,{\\rm GeV}$ detected by EGRET (\\cite{fic96,sre98}) as coming from quiescent-state unresolved blazars. Such a background has not yet been detected in the VHE range, as it is technically difficult with the atmospheric Cherenkov technique to find an isotropic extragalactic emission and even more to distinguish it from the cosmic-ray electron flux (\\cite{egb08}). In addition, the EBL attenuation limits the distance from which $\\sim$TeV $\\gamma$-rays can propagate to $\\sim 1\\,{\\rm Gpc}$ (\\cite{HESSEBL}). As pointed out by \\cite{che00}, emission mechanisms might be simpler to understand during quiescent states in blazars, and they are also the most likely state to be found observationally. In the X-ray band, the existence of a steady underlying emission has also been invoked for two other VHE emitting blazars (Mrk~421, \\cite{fos00}, and 1ES~1959+650, \\cite{gie02}). Being able to separate, and detect, flaring and nonflaring states in VHE $\\gamma$-rays is hence important for such studies.  The observation of the spectacular outbursts of \\pks\\ in July 2006 represents one the most extreme examples of AGN variability in the TeV domain, and allows spectral and timing properties to be probed over two orders of magnitude in flux.  Whereas for the flaring states with fluxes above a few $10^{-10}\\,\\percmsqrs$ a clear hardening of the spectrum with increasing flux is observed, familiar also from the blazars Mrk 421 and Mrk 501, for the quiescent state in contrast an indication of a softening is noted. If confirmed, this is a new and intruiging observation in the VHE regime of blazars. The blazar PKS~0208$-$512 (of the FSRQ class) also shows such initial softening and subsequent hardening with flux in the MeV range, but no general trend could be found for $\\gamma$-ray blazars (\\cite{nan07}). In the framework of synchrotron self-Compton scenarios, VHE spectral softening with increasing flux can be associated with, for example, an increase in magnetic field intensity, emission region size, or the power law index of the underlying electron distribution, keeping all other parameters constant. A spectral hardening can equally be obtained by increasing the maximal Lorentz factor of the electron distribution or the Doppler factor (see e.g. Fig.~11.7 in \\cite{kat99}). A better understanding of the mechanisms at play would require multi-wavelength observations of similar time span and sampling density as the data set presented here.  It is shown that the variability time scale $t_{\\rm var}$ of a few minutes are only upper limits for the intrinsic lowest characteristic time scale. Doppler factors of $\\delta \\geq 100$ of the emission region are derived by \\cite{HESS2155_BigFlare}~(2007a) using the $\\sim$~$10^9~M_\\odot$ black hole (BH) Schwarzschild radius light crossing time as a limit, while \\cite{begelman} argue that such fast time scales cannot be linked to the size of the BH and must occur in regions of smaller scales, such as ``needles'' of matter moving faster than average within a larger jet (\\cite{ghi08}), small components in the jet dominating at TeV energies (\\cite{kat08}), or jet ``stratification'' (\\cite{bou08}). \\cite{levinson} attributes the variability to dissipation in the jet coming from radiative deceleration of shells with high Lorentz factors.  The flaring period allowed the study of light curves in separated energy bands and the derivation of a power law dependence of $F_{\\rm var}$ with the energy ($F_{\\rm var}\\propto E^{\\sim 0.2}$). This dependence is comparable to that reported in \\cite{berrie}~(2007), \\cite{integr_mrk}~(2008), \\cite{xmm}~(2002), where $F_{\\rm var}(E)\\propto E^{\\sim0.2}$ between the optical and X-ray energy bands was found for Mrk\\,421 and \\pks, respectively. An increase with the energy of the flux variability has been found for Mrk\\,501 (\\cite{magic_mrk501}~2007) in VHE $\\gamma$-rays on timescales comparable to those observed here.   The flaring period showed for the first time that the intrinsic variability of \\pks\\ increases with the flux, which can itself be described by a lognormal process, indicating that the aperiodic variability of \\pks\\ could be produced by a multiplicative process. The flux in the ``quiescent regime'', which is on average 50 times lower than in the flaring period and has a 3 times lower $F_{\\rm var}$, also follows a lognormal distribution, suggesting similarities between these two regimes.  It has been possible to characterize a power spectral density of the flaring period in the frequency range $10^{-4}$--$10^{-2}\\,{\\rm Hz}$, resulting in a power law of index $\\alpha=2.06\\pm0.21$ valid for frequencies down to $\\sim 1/{\\rm day}$. The description of the rapid variability of a TeV blazar as a random stationary process must be taken into account by time-dependent blazar models. For \\pks~the evidence of this log-normality has been found very recently in X-rays (\\cite{pks_xrays}) and  as previously mentioned, X-ray binaries and Seyfert galaxies also show lognormal variability, which is thought to originate from the accretion disk (\\cite{mchardy04,lyubarskii,arevalo}), suggesting a connection between the disk and the jet. This variability behaviour should therefore be searched for in existing blazar light curves, independently of the observed wavelength."
    },
    "1005/1005.1707_arXiv.txt": {
        "abstract": "The evolution of primordial collapsing clouds and formation of proto-Population III stars are investigated using three-dimensional ideal MHD simulations. We calculated the collapse of magnetized primordial clouds from the prestellar stage until the epoch after the proto-Population III star formation, spatially resolving both parsec-scale clouds and sub-AU scale protostars. The formation process of proto-population III star is characterized by the ratio of rotational to magnetic energy of the natal cloud. When the rotational energy is larger than the magnetic energy, fragmentation occurs in the collapsing primordial cloud before the proto-Population III star formation and binary or multiple system appears. Instead, when the magnetic energy is larger than the rotational energy, strong jet with $>100\\,{\\rm km\\,s}^{-1}$ is driven by circumstellar disk around the proto-population III star without fragmentation. Thus, even in the early universe, the magnetic field plays an important role in the star formation process. ",
        "introduction": "Magnetic fields is a key ingredient in present-day star formation. For example, protostellar jets, which are ubiquitous in star-forming regions, are considered to be driven from protostars by the Lorentz force. Protostellar jets influence gas accretion onto protostars and disturb the ambient medium. In addition, the angular momentum of the cloud is removed by magnetic braking and protostellar jets. The removal of angular momentum makes protostar formation possible in a parent cloud that has a much larger specific angular momentum than the protostar. So far, magnetic effects in primordial gas clouds have been ignored in many studies because magnetic fields in the early universe are supposed to be extremely weak. However, recent studies indicate magnetic fields of moderate strength can exist even in the early universe. Cosmological fluctuations produced magnetic fields before the epoch of recombination \\citep{ichiki06}. These fields were sufficiently large to seed the magnetic fields in galaxies. A generation mechanism for magnetic fields at the epoch of reionization are also proposed \\citep{langer03}, in which  magnetic fields in intergalactic matter are amplified up to $\\sim 10^{-11}\\ {\\rm G}$. These fields can therefore increase up to $\\sim 10^{-7}-10^{-8}\\ {\\rm G}$ in the first collapsed object having number density of $n \\sim 10^3 \\cm$. These fields may influence the evolution of  primordial gas clouds and formation of Population III stars.  Under spherical symmetry including  hydrodynamical radiative transfer, many authors have carefully investigated both the present-day \\citep[e.g.,][]{masunaga00} and primordial \\citep[e.g.,][]{omukai98} star formation processes. A significant difference between present-day and primordial star formation exists in the thermal evolution of the collapsing gas cloud because of differences in the abundance of dust grains and metals. In present-day star formation, the gas temperature in molecular clouds is $\\sim10$\\,K. These clouds collapse isothermally for $\\nc \\lesssim 10^{11}\\cm$. Then the gas becomes adiabatic at $\\nc \\simeq 10^{11}\\cm$, and an adiabatic core (or the first core) forms. After the dissociation of molecular hydrogen ($\\nc \\simeq 10^{16}\\cm$), the protostar forms at $\\nc \\simeq 10^{21}\\cm$. On the other hand, primordial gas clouds have temperatures of $\\sim200-300$\\,K at $\\nc \\simeq 10^3\\cm$ \\citep{omukai05,bromm02}. These clouds collapse keeping polytropic index $\\gamma\\simeq 1.1$ for a long range of $10^4\\cm \\lesssim \\nc \\lesssim 10^{16}\\cm$. Thus, the first core does not appear in the primordial collapsing cloud. After the central density reaches $\\nc \\simeq 10^{16}\\cm$, the thermal evolution of the primordial collapsing cloud begins to coincide with that of a present-day cloud  \\citep{omukai05}. The difference in thermal evolution between present-day and primordial clouds arises when $\\nc \\lesssim 10^{16}\\cm$.   Another major difference between present-day and primordial star formation exists in their magnetic evolution. In present-day star formation, neutral gas is well-coupled  with ions for $\\nc \\lesssim 10^{12}\\cm$ and $\\nc \\gtrsim 10^{15}\\cm$, while the magnetic field dissipates by Ohmic dissipation in the range of $10^{12}\\cm \\lesssim \\nc \\lesssim 10^{15}\\cm$ \\citep{nakano02}. In the collapsing cloud, $\\sim99$\\% of the magnetic field  is dissipated for $10^{12}\\cm \\lesssim \\nc \\lesssim 10^{15}\\cm$ \\citep{machida07}. On the other hand, in a primordial gas cloud,  the magnetic field couples strongly with the primordial gas during all phases of the star formation, as long as the initial field strength is weaker than $B_0\\lesssim 10^{-5}(n/10^3\\cm)^{0.55}$\\,G \\citep{maki04,maki07}. In summary, the magnetic field is largely dissipated by Ohmic dissipation before protostar formation in present-day clouds, while the magnetic field can continue to be amplified without dissipation in primordial clouds.  The magnetic field in the collapsing cloud is closely related to the fragmentation or formation of binary or multiple stellar systems. In present-day star formation, the magnetic field strongly suppresses rotation-driven fragmentation \\citep{machida08b}. For primordial clouds, magnetic effects on fragmentation are still unknown. In this study, we investigate the evolution of weakly magnetized primordial clouds and the formation of Population III stars using three-dimensional simulations. ",
        "conclusions": "In this study, we calculated cloud evolution from the stage of $n_c = 10^3\\cm$ until the protostar is formed ($\\simeq 10^{22}\\cm$) for 36 models, parameterizing the initial magnetic field strength and rotation, to investigate effects of magnetic fields in collapsing primordial clouds. Our calculations showed that fragmentation occurs but no jet appears when $\\beta_0 > \\gamma_0$, and jet appears after the protostar formation without fragmentation when $\\beta_0 < \\gamma_0$. Thus, in the collapsing primordial cloud, the cloud evolution is mainly controlled by the centrifugal force than the Lorentz force when $\\beta_0 > \\gamma_0$, while the Lorenz force is more dominant than the centrifugal force when $\\gamma_0  > \\beta_0$.  A jet is driven when the initial cloud has magnetic field of \\begin{equation} B_0 \\gtrsim 10^{-9} \\left( \\dfrac{\\nc}{10^3\\cm} \\right)^{2/3} \\,G \\label{eq:bcrit} \\end{equation} if the cloud rotates slowly as $\\Omega \\lesssim 4\\times 10^{-17}(\\nc/10^3\\cm)^{2/3}$\\,s$^{-1}$. The power of a jet, e.g., a mass ejection rate, is considered to be controlled by the accretion rate as indicated in present-day star formation; the mass ejection rate of a jet is 1/10 of the mass accretion onto the protostar. The accretion rate of primordial star formation is expected to be considerably larger than that at present day, and it produces a stronger jet. The life time of the jet also seems to be controlled by accretion in the present-day; a jet stops when mass accretion stops. For Population III stars, the gas accretion does not halt within their lifetimes \\citep{omukai01,omukai03}. Therefore, a jet also may continue during the all lifetime of the protostar, and the strong jet propagates to disturb a surrounding medium significantly. The disturbance of the medium could trigger the subsequent star formation as frequently observed in present-day star formation.  Assuming the power law growth of $B_0 \\propto n_c^{2/3}$, the critical strength of the magnetic field, $B_0 = 10^{-9}$\\,G at $n_c = 10^{3}\\cm$ corresponds to $B_0 = 5\\times 10^{-13}$\\,G at $n_c = 0.01\\cm$ [see, eq.~(\\ref{eq:bcrit})], which is much stronger than the background magnetic field  derived by \\citet{ichiki06}. However, when the magnetic field is amplified to $B \\sim 10^{-9}(\\nc/10^3\\cm)^{2/3}$\\,G by some mechanisms \\citep[e.g.,][]{schleicher10}, the magnetic field can affect the collapse of the primordial cloud. Even if a cloud has a magnetic field weaker than the critical strength $B_0 = 10^{-9}$\\,G, the magnetic field may play an important role after the protostar formation. In analytical study of the evolution of accretion disks around the first stars \\citep{tan04} , it is suggested  that magnetic fields amplified in the circumstellar disk eventually give rise to protostellar jets during the protostellar accretion phase.  Rotation promotes fragmentation when the first collapsed objects has the angular velocity of $\\Omega_0 \\gtrsim 10^{-17}(\\nc/10^3\\cm)^{2/3}$\\,s$^{-1}$. The fragmentation is expected to produces binary or multiple stellar system. When a multiple stellar system is formed, some stars can be ejected by close encounters. At the protostar formation epoch, the protostar has a mass of $M\\simeq 10^{-3}\\msun$. The ejected proto-Population III stars may evolve to metal-free brown dwarfs or low-mass stars. When a binary component in a multiple stellar system is ejected from the parent cloud by protostellar interaction, a low-mass metal free binary may also appear in the early universe. It is considered that the extremely metal-poor ([Fe/H]$<$-5) stars \\citep{christlieb01,frebel05} are formed as binary members from metal-free gas, and then have been polluted by the companion stars during the stellar evolution \\citep{suda04}. In addition, a binary frequency in Population III star may be comparable to or larger than that at present day \\citep{komiya06,machida08c,machida08d,machida09a,machida09b}. In order to confirm the ejection scenario, the further long-term calculations are required."
    },
    "1005/1005.1641_arXiv.txt": {
        "abstract": "In the context of the VLA-COSMOS Deep project additional VLA A array observations at 1.4\\,GHz were obtained for the central degree of the COSMOS field and combined with the existing data from the VLA-COSMOS Large project. A newly constructed Deep mosaic with a resolution of 2.5$''$ was used to search for sources down to 4$\\sigma$ with 1$\\sigma \\approx 12\\,\\mu$Jy/beam in the central 50$'\\times$50$'$. This new catalog is combined with the catalog from the Large project (obtained at 1.5$''$$\\times$1.4$''$ resolution) to construct a new Joint catalog. All sources listed in the new Joint catalog have peak flux densities of $\\ge$5$\\sigma$ at 1.5$''$ and/or 2.5$''$ resolution to account for the fact that a significant fraction of sources at these low flux levels are expected to be slighty resolved at 1.5$''$ resolution. All properties listed in the Joint catalog such as peak flux density, integrated flux density and source size are determined in the 2.5$''$ resolution Deep image. In addition, the Joint catalog contains 43 newly identified multi-component sources. ",
        "introduction": "In recent years, several cosmological deep fields have been imaged at 20\\,cm \\citep[e.g., ][]{ric00,bon03,con03,hop03,sey04,nor05,huy05,fom06,sim06,ivi07,sch07,mil08,owe08} providing a few thousand radio sources down to flux limits of a few 10\\,$\\mu$Jy. These deep radio imaging data are sensitive enough to detect star forming galaxies with star formation rates of several 10 to 100 $\\solmy$ out to and beyond a redshift of $z$\\,$\\sim$\\,1. Similarly, radio galaxies can be seen out to redshifts of $z$\\,$\\sim$\\,5 and the most luminous ones even well into the epoch of reionization. Thus, deep radio images in conjunction with deep imaging data at X-ray, optical and infrared wavelengths are ideal to investigate the dust-unbiased star formation, the evolution of radio(-loud) AGN, as well as the population mix of radio sources in the first place.  In order to study the cosmological evolution of galaxies and black holes, it is not only important to overcome the effect of cosmic variance (e.g. by studying a large enough area) but also to understand the effect of large scale structure on the evolution (e.g. by covering a large contiguous area). To address the second effect in particular, the Cosmic Evolution Survey (COSMOS)\\footnote{~\\tt http://cosmos.astro.caltech.edu} collaboration has conducted panchromatic imaging and spectroscopy of an equatorial field with a size of 2\\,$\\rm deg^2$ \\citep[for an overview, see ][]{sco07a} ranging from X-ray XMM-Newton and Chandra \\citep{has07,elv09}, UV GALEX \\citep{zam07}, optical and near-infrared ground-based \\citep{tan07,cap07a}, optical HST \\citep{sco07b,koe07}, mid- to far-infrared Spitzer \\citep{san07}, millimeter \\citep{ber07,sco08} and radio VLA \\citep{sch04,sch07} imaging to extensive optical spectroscopy using the VLT/VIMOS and Magellan/IMACS instruments \\citep{lil07,tru07}. Most of these datasets are now publicly available from the COSMOS archive at IPAC/IRSA\\footnote{~\\tt http://irsa.ipac.caltech.edu/Missions/cosmos.html}.  The VLA-COSMOS survey at 20\\,cm is part of the overall imaging effort and its scientific goals and motivation have been described in detail by \\cite{sch07}. Initial observations from a pilot project testing the mosaicking strategy and giving a first source catalog are presented by \\cite{sch04}. As a large NRAO/VLA program, the VLA was used in A and C configuration to cover the entire COSMOS field resulting in an image with uniform noise properties in the central 1$\\times$1\\,deg$^2$ and an average rms of 10.5\\,$\\mu$Jy. \\cite{sch07} provide a detailed description of the survey set-up, the data reduction, as well as the testing and construction of the final VLA-COSMOS Large project catalog (hereafter: Large catalog). Subsequently, \\cite{bon08} derived the completeness of the Large catalog and also analyzed the effect of bandwidth smearing on the derived source flux densities to obtain the source counts. Although the VLA-COSMOS Large Project dataset has been used for several scientific results on, e.g. the faint radio population, the radio-derived star formation rate density, the radio AGN population and stacking of high-z galaxy populations \\citep{smo08,smo09a,smo09b,car07,car08}, the need for deeper radio imaging data became apparent during the search for radio counterparts to millimeter sources from the COSMOS MAMBO mapping data \\citep{ber07}; the Large project only provided counterparts for about half of the mm-sources. Thus, the Deep project was initiated with the aim of doubling the integration time for the central seven pointings which fully cover the MAMBO 1.2\\,mm map of the COSMOS field. These new observations are described here.  The paper is organized as follows: After a description of the new observations and the data reduction (\\S \\ref{sec:observations}), the revision of the Large catalog (leading to a new version of v2.0) and the construction of the new Deep catalog are outlined in \\S \\ref{sec:newLarge} and \\S \\ref{sec:deepcat}. In \\S \\ref{sec:identify} we explain the construction of the final Joint catalog which is described in detail in \\S \\ref{sec:joint_cat} where we also present all the tests and corrections involved. A summary is given in \\S \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  Continued analysis of the VLA-COSMOS catalog presented in \\cite{sch07} and the completion of the VLA-COSMOS Deep project motivated the compilation of a new radio catalog for the COSMOS field. The VLA-COSMOS Joint catalog was generated by combining the catalogs of the VLA-COSMOS Large Project with a newly created source catalog (Deep catalog) from the 2.5$''$ resolution Deep mosaic. This catalog is already available for download by the public at the COSMOS archive at IPAC/IRSA\\footnote{~~\\texttt{http://irsa.ipac.caltech.edu/data/COSMOS/tables/vla}}. A comparison of the depth and areal coverage for a representative sample of deep field radio surveys at 1.4\\,GHz shows that the VLA-COSMOS covers the largest area at its depth and angular resolution (Fig. \\ref{fig:radio_surveys}). Thus it should be well suited to also study effects of the Large Scale Structure on the presence/absence of radio emission.  The reduction and analysis of the deeper 20\\,cm observations of the central 7 pointings of the VLA-COSMOS projects using the VLA in A configuration have been described in detail (also referred to as VLA-COSMOS Deep project). In order to minimize the effect of bandwidth smearing the Deep mosaic has a resolution of 2.5$''$ (compared to 1.5$''$ for the Large mosaic) and it was used to create a corresponding source catalog using the task \\texttt{SAD}.  An input list for the new Joint catalog was compiled by combining the revised Large catalog (v2.0) and the new Deep catalog. The criteria were set such that no particular bias against slightly extended radio sources was present when selecting the sources. All properties of the radio sources listed in the Joint catalog have been derived in the 2.5$''$ resolution Deep mosaic.  The construction of the Joint catalog was motivated by the desire to provide a catalog of bona-fide radio sources in the COSMOS field for distinct science applications that are interested in the radio properties of certain populations of galaxies. On the other hand the revised Large catalog (v2.0) is flux-limited (in radio), has a fairly uniform sensitivity coverage and its completeness is well characterized \\citep[see][]{bon08}, thus it is well suited for, e.g., studies of the faint radio population \\citep[such as, e.g.,][]{smo08,smo09a,smo09b}.  \\appendix"
    },
    "1005/1005.4855_arXiv.txt": {
        "abstract": "In this paper we present in detail the methodology and the first results of a ground-based program to determine the absolute proper motion of the Fornax dwarf spheroidal galaxy.   The proper motion was determined using bona-fide Fornax star members measured with respect to a fiducial at-rest background spectroscopically confirmed Quasar, \\qso. Our homogeneous measurements, based on this one Quasar gives a value of (\\mua,\\mud)$ = (0.64 \\pm 0.08, -0.01 \\pm 0.11)$~\\masy. There are only two other (astrometric) determinations for the transverse motion of Fornax: one based on a combination of plates and HST data, and another (of higher internal precision) based on HST data. We show that our proper motion errors are similar to those derived from HST measurements on individual QSOs. We provide evidence that, as far as we can determine it, our motion is not affected by magnitude, color, or other potential systematic effects. Last epoch measurements and reductions are underway for other four Quasar fields of this galaxy, which, when combined, should yield proper motions with a weighted mean error of $\\sim50\\,\\mu$as~y$^{-1}$, allowing us to place important constraints on the orbit of Fornax. ",
        "introduction": "The proper motions (PMs) of the satellites of the Milky Way (MW), when combined with existing radial velocities, allow us to determine the space velocity vectors of these satellites of our galaxy, which in turn place important constraints on their orbits (see, e.g., \\citet{bes07}, \\citet{pia07}). This knowledge is crucial to determine if these galaxies are gravitationally bound to the Galaxy, and to our understanding of the evolution and origin of its satellite system (\\citet{byr94}). The PMs of the satellites of the MW are necessary to understand: a) the origin of the MW satellite system and its relationship with the formation and evolution of the galactic halo (\\citet{din00}, \\citet{pal02}, \\citet{car09}), b) the nature and origin of the streams that seem to align different subgroups of these galaxies (\\citet{lyn82}, \\citet{lyn95}, \\citet{pia05}), and c) the role of tidal interactions in the evolution and star formation history of low mass galaxies (\\citet{zar04}, \\citet{pia05}, \\citet{noe09}, \\citet{may10}). A comprehensive study of the MW's satellite system can lead to a greater general understanding of galaxy evolution and the physical processes governing star formation in galaxies.   From another perspective, reliable space motions of Local Group galaxies is a key ingredient to populate the phase-space components for flow models which predict the dynamical evolution of the local universe. Indeed, current center-of-mass locations and motions from the present distribution of galaxies can be used as boundary conditions to ``trace them back in time'' \\citep{pee89,pee94} and to test the paradigm that galaxy clusters grew by gravitational instabilities from an originally smoother medium, with small random motions. These earlier simulations have not been repeated too often due to the growing realization that more complex effects such as tides, satellite interactions, accretion and mergers, influencing the growth of galaxies with time (specially at epochs earlier than 7-8~Gyr from now), can greatly complicate this approach. Nevertheless, the motions of local group galaxies can be used as boundary conditions in a first aproximation to carry out the above analysis. This goal has become one of the important drivers behind future astrometric space missions, such as SIM (\\citet{unw08}, specially their Figure~13), which expect to measure PM for Local Group galaxies with a precision ten times better than what can be achieved with present techniques.  With the above motivations in mind, in the year 2000 we started a ground-based program aimed at determining, the absolute PM of three southern dwarf Spheroidal (dSph) galaxies, Carina, Fornax, and Sculptor, with respect to known background Quasars (QSOs) that can be used as inertial reference points. Three epochs, over a period of eight years were obtained using a single telescope+detector set up: Ours is then the first entirely optical CCD/ground-based proper motion study of an external galaxy other than the Magellanic Clouds.  In this paper we report on the first results from this program, based on one QSO field in Fornax, for which we have data of good enough quality to allow us to asses the expected precision of our measurements, and to describe our methodology in detail. Last epoch measurements and reductions are underway for other four Quasar fields of Fornax, as well as for a similar number of QSO fields in Carina and Sculptor. The results for these will be presented in forthcoming papers.  In section~\\ref{obma} we describe our observational material and data acquisition strategies, in section~\\ref{rest} we describe our methodology for deriving the PMs, in section~\\ref{anal} we present the analysis for our QSO field, and in section~\\ref{comp} we present our main conclusions as well as a comparison to other results.  \\newpage ",
        "conclusions": "\\label{comp}  There have been only two astrometric determinations of the PM for the Fornax dSph galaxy, namely that by \\citet{din04}, based on a combination of ground-based plates and Hubble-WFPC data, and that based exclusively on HST data (\\citet{pia07}, which gives revised values to those reported earlier in \\citet{pia02}).  \\citet{din04}, based on an independent determination using 48 galaxies and 8 QSOs give a weighted mean of (\\mua,\\mud)=($0.59 \\pm 0.16, -0.15 \\pm 0.16$)~\\masy. On the other hand, \\citet{pia07}, based on 4 QSOs (their Table~3), gives a weighted mean of (\\mua,\\mud)=($0.476 \\pm 0.046, -0.360 \\pm 0.041$)~\\masy. In order to compare our result (based on a single QSO) with their measurements it makes more sense to compare our value with their {\\it individual} PM values. This allows us to estimate the expected uncertainty on our final weighted mean when we incorporate our other 4 QSO fields. In Figure~\\ref{pmcomp} we plot the individual measurements from these previous works, along with our measurement. We note however that, while \\citet{pia07} report individual measurements, \\citet{din04} only give the mean values with respect to 8 QSOs and 48 galaxies, and these are the results plotted in Figure~\\ref{pmcomp}. \\citet{din04} {\\it do} give the {\\it individual} PMs for images A and B of the same QSO reported in this paper. These individual values are also plotted in Figure~\\ref{pmcomp}. By looking at the open and filled squares we can see the improvement achieved by ground-based astrometry when using CCDs and a homogeneous data set in comparison with the results that combine non-linear plates and heterogeneous data. Nevertheless, it is somewhat surprising how close our (single) value is to the mean PM from galaxies as derived by \\cite{din04} (differences smaller than $1\\sigma$ of {\\it our} smaller error), while there is a larger discrepancy (although still smaller than $2\\sigma$ of {\\it their} (larger) error) when compared to the PM using QSOs. In Dinescu's study they have a large sample of brighter better-measured galaxies (see their Figure~2) while the (much fewer) QSOs are fainter and possibly have larger positional errors (although, their individual PM errors are smaller than those of galaxies for a given magnitude all the way down to V$\\sim 21$, possibly due to the extended, more diffuse nature of galaxies). Overall, their mean PM with respect to galaxies has an error $\\sim50$\\% smaller than that derived from the QSOs.  In comparison with the individual HST measurements, we see a rather large discrepancy, specially in DEC: While our value lies between 2.3 and 1.2$\\sigma$ away from HST individual PMs in \\mua, this difference becomes between 4.4 and 2.4$\\sigma$ in \\mud. We note that the largest \\& smallest PM values from HST for {\\it both} \\mua\\ and \\mud\\ (see Table~3 on \\citet{pia07}) come precisely from images A and B of \\qso. \\citet{din04} also point the rather large discrepancy in the PMs derived from components A and B (plotted in our Figure~\\ref{pmcomp}). As was described in section~\\ref{centroids} the A component of QJ~0240-3434 {\\it is} indeed affected by two close companions (see Figure~\\ref{qsoab}). It is not unlikely that tiny changes in the photocenter of the images due to other close, unresolved and fainter, companions, or even slightly extended structure(s) in the wings of the QSO (from the underlying QSO galactic disk) could introduce an extra source of noise in the PMs that is not necessarily accounted for in the final PM error budget. Obviously, whether or not our own measurements are subject to some (as yet unknown) systematic effect related to what we just mentioned, is an issue that could be addressed once we incorporate the other QSO fields. In Table~\\ref{pmtab} we summarize the measurements plotted in Figure~\\ref{pmcomp} (the ``perspective'' motions are explained further below).  It is interesting to ascertain to what type of final (Heliocentric) velocity errors correspond our individual PM measurement errors. For a given distance and distance error $r \\pm \\sigma_r$ and PM and its error in component ``x'' (in this case either RA or DEC), $\\mu_x \\pm \\sigma_{\\mu_x}$, the corresponding velocity and its error is given by:  \\begin{eqnarray} v_x = K \\, r \\, \\mu_x \\label{vel} \\\\ \\nonumber \\\\ \\frac{\\sigma_{v_x}}{v_x} = \\sqrt{ \\left( \\frac{\\sigma_r}{r} \\right)^2 + \\left( \\frac{\\sigma_{\\mu_x}}{\\mu_x} \\right)^2 } \\label{velerr} \\end{eqnarray}  where, if $r$ is in kpc, and $\\mu_x$ is in \\masy, then $K=4.74$ and $v_x$ and $\\sigma_{v_x}$ are in km~s$^{-1}$. According to \\citet{mat98}, the distance to Fornax is $138 \\pm 8$~kpc (i.e., 6\\% formal error). Therefore in equation~(\\ref{velerr}) the distance error is negligible in comparison with our PM errors in either RA or DEC. For our PM reference value (see Table~\\ref{pmtab}) we obtain a Heliocentric velocity and error of $v_{\\alpha}=419 \\pm 58$~km~s$^{-1}$ and $v_{\\delta}= -7 \\pm 72$~km~s$^{-1}$. The measured Heliocentric radial velocity from \\citep{mat98} is $v_r = 53 \\pm 3$~km~s$^{-1}$. It is thus clear that the derived Heliocentric Fornax motion is dominated by uncertainties in the PM from our one-QSO measurement, not by the distance or radial velocity uncertainty. Given our measurement errors, once we incorporate our other four QSO fields, we expect to achieve velocity errors for the weighted mean on the order of 30~km~s$^{-1}$ per component, totally compatible with the (weighted mean) HST results. Table~\\ref{pmtab} also shows that our proper motion value exhibits the largest tangential velocity for Fornax after the Dinescu et. al value from Galaxies. At this stage we howevere refrain from making a full analysis of the implications of this result, until we have collected the data from the other four QSO fields which will allow us to have a more robust weighted mean.  We must note that, when comparing PMs values derived for an extended object, such as Fornax, one must actually compare ``center-of-mass'' (COM) motions to avoid any projection effects and internal galaxy motions (e.g., galactic rotation) that might alter the observed motions. For computing galactic orbits of these galaxies we also need COM velocities. Our measurements are performed, however, on fields away from the COM, and we have to apply corrections to account for this situation. As explained in \\citet{cos09}, in the case of the LMC our PMs had to be corrected not only for projection effects (see below), but also by the effect on our measured PMs induced by the (differential) rotation of the plane of the LMC. To apply these corrections, we followed the prescriptions described by \\cite{jon94} and \\cite{van02}, which assumed that our fields lay in the plane of the LMC (or SMC) and shared the motion of its disk. The situation for the dSphs is quite different however, they do not exhibit any hint of large-scale rotation, nor of the existence of a disk or other well-defined structure; rather they show a smooth spheroidal distribution and the motion of its stars is dominated by their velocity dispersion, which is typically a few~km~s$^{-1}$ \\citep{van99}. This velocity dispersion translates into a PM dispersion of a few $\\mu$as~y$^{-1}$ which is not measurable by current astrometric techniques (although it will be measured in the future by SIM). Therefore, if we assume that there are no large-scale streaming motions in Fornax, our PMs are not affected, as far as we can measure it, by internal kinematic effects. We do, however, need to correct for purely geometrical projection effects, this is done as described in the following paragraph.  If $(v_r,v_\\alpha,v_\\delta)$ are the observed (measured) Heliocentric radial velocity, velocity in RA and velocity in DEC respectively, for a field at position $(\\alpha,\\delta)$ and Heliocentric distance $r$, the velocity for the COM $(v^*_r,v^*_\\alpha,v^*_\\delta)$ at position $(\\alpha^*,\\delta^*)$ and distance $r^*$ is given by:  \\begin{equation} \\left( \\begin{array}{c} v^*_r \\\\ v^*_\\alpha \\\\ v^*_\\delta \\end{array} \\right) = M^{-1} (\\alpha^*,\\delta^*) \\cdot M(\\alpha,\\delta) \\cdot \\left( \\begin{array}{c} v_r \\\\ v_\\alpha \\\\ v_\\delta \\end{array} \\right) \\label{proj} \\end{equation}  where $M(\\alpha,\\delta)$ is a rotation matrix whose components are:  \\begin{equation} M(\\alpha,\\delta) = \\left( \\begin{array}{ccc} \\cos \\delta \\cos \\alpha & -\\sin \\alpha  & -\\sin \\delta \\cos \\alpha \\\\ \\cos \\delta \\sin \\alpha &   \\cos \\alpha & -\\sin \\delta \\sin \\alpha \\\\ \\sin \\delta             & 0              &  \\cos \\delta \\end{array} \\right) \\label{matrix} \\end{equation}  and which satisfies that $M^{-1}(\\alpha,\\delta)=M^t(\\alpha,\\delta)$. While $r$ and $r^*$ are not explicitely written in equation~(\\ref{proj}), they are implicitly used to go from PMs to tangential velocities (or vice versa) through equation~(\\ref{vel}). Of course, $r^* = 138 \\pm 8$~kpc. In the case of the LMC (and SMC) \\citet{cos09} assumed that our fields were located in a disk-like structure with known orientation in the sky (inclination and line of nodes). We do not have however such structures in the featureless dSphs. If we assume, e.g., as a first crude approximation that the galaxy lies in a plane in the sky perpendicular to the line-of-sight to the COM. In this case one can show that $r$ is given by:  \\begin{equation} r = \\frac{r^*}{\\cos \\delta \\cos \\alpha \\cos \\delta^* \\cos \\alpha^* + \\cos \\delta \\sin \\alpha \\cos \\delta^* \\sin \\alpha^* + \\sin \\delta \\sin \\delta^*} \\end{equation}  We note that, in equation~(\\ref{proj}) $v_r$ is {\\it not} known, but we do know $v^*_r = 53 \\pm 3$~km~s$^{-1}$ from the literature. Therefore, starting from our measured $v_\\alpha$ and $v_\\delta$ values we iterate on the $v_r$ values until we reproduce the expected $v^*_r$ for the COM, in a procedure similar to that adopted by \\citet{cos09}. Also, equations (\\ref{proj}) and (\\ref{matrix}) easily allow us to fully propagate errors on all the measured quantities (radial velocities, proper motions, distances) from the observed values to the sought-for COM values. For our measured PM, the COM distance and radial velocity indicated previously, and the Fornax COM position given by \\citet{mat98}, we obtain $v^*_{\\alpha}=419 \\pm 52$~km~s$^{-1}$ and $v^*_{\\delta}= -7 \\pm 72$~km~s$^{-1}$, i.e., a value similar to that computed without {\\it any} projection correction. This is due to the fact that, in this particular case, these corrections are actually quite small, the PM corrections are 0.05~$\\mu$as in RA and by 0.45~$\\mu$as in DEC. These corrections are, of course, much smaller than the measurement uncertainties involved (this was not case of the LMC and SMC fields reported in \\citet{cos09}). Note that the field reported here is quite close to the Fornax COM, for some of the other more distant QSO fields in our program, these corrections might be slightly larger, but can be readily computed in each case. None of the values reported by other authors in Table~\\ref{pmtab} have been corrected for any projection effect, but since these corrections are tiny, one can readily compare them without further corrections.  For completeness, in Table~\\ref{pmtab} (and in Figure~\\ref{pmcomp}), we have included the recent determination of the Fornax PM by \\citet{wal08}, using what they call the ``perspective rotation'' method (in a way, this method is the reverse of the ``astrometric radial velocities'', described by, e.g., \\citet{lin00}). As it can be easily seen from Equations~\\ref{proj} and~\\ref{matrix}, one could write an equivalent equation for $(v_r, v_\\alpha, v_\\delta)$ as a function of $(v^*_r, v^*_\\alpha, v^*_\\delta)$ by a simple matrix inversion. The first row of that equation would give the observed radial velocity $v_r$ as a function of a combination of $(v^*_r, v^*_\\alpha, v^*_\\delta)$ (which, for a given galaxy is of course a fixed quantity - independent of the field observed). Therefore, if one measures the $v_r$ of samples of stars at different locations $(\\alpha,\\delta)$ across a galaxy, one can solve through some minimization algorithm for the unknown $(v^*_r, v^*_\\alpha, v^*_\\delta)$. \\citet{wal08} have used precisely this approach to determine the ``perspective'' PMs for 4 dSphs, including Fornax, using radial velocities exclusively. As it can be seen from Table~\\ref{pmtab}, the method has errors compatible with the best purely astrometric determinations, and can thus become a useful complement to them.  \\newpage"
    },
    "1005/1005.3472_arXiv.txt": {
        "abstract": "{We present far-infrared spectroscopic observations, taken with the Photodetector Array Camera and Spectrometer (PACS) on the \\emph{Herschel Space Observatory}, of the protoplanetary disk around the pre--main--sequence star HD\\,100546. These observations are the first within the DIGIT \\emph{Herschel} key program, which aims to follow the evolution of dust, ice, and gas from young stellar objects still embedded in their parental molecular cloud core, through the final pre--main--sequence phases when the circumstellar disks are dissipated.} {Our aim is to improve the constraints on temperature and chemical composition of the crystalline olivines in the disk of HD\\,100546 and to give an inventory of the gas lines present in its far-infrared spectrum. } {The 69\\,$\\mu$m feature is analyzed in terms of position and shape to derive the dust temperature and composition. Furthermore, we detected 32 emission lines from five gaseous species and measured their line fluxes. } {The 69\\,$\\mu$m emission comes either from dust grains with $\\sim$70~K at radii larger than 50\\,AU, as suggested by blackbody fitting, or it arises from $\\sim$200\\,K dust at $\\sim$13\\,AU, close to the midplane, as supported by radiative transfer models. We also conclude that the forsterite crystals have few defects and contain at most a few percent iron by mass. Forbidden line emission from [C~{\\sc ii}] at 157\\,$\\mu$m and [O~{\\sc i}] at 63 and 145\\,$\\mu$m, most likely due to photodissociation by stellar photons, is detected.  Furthermore, five H$_2$O and several OH lines are detected. We also found high-J rotational transition lines of CO, with rotational temperatures of $\\sim$300\\,K for the transitions up to $J=22-21$ and $T\\,\\sim$800\\,K for higher transitions.  } {} ",
        "introduction": "Circumstellar disks around young stars are the birthplaces of planetary systems. To understand planet formation, it is vital to study the processes that govern the evolution of gas and dust in these disks. PACS provides unique information in this field through observations of far-infrared (IR) solid-state features which are particularly sensitive to temperature and elemental composition. Moreover, PACS is well suited to study the warm gas of a few 100\\,K, complementary to the hot gas probed in near-IR and the cold gas observed at millimeter wavelengths.   An intensively studied pre-main-sequence star is the Herbig B9.5Vne star HD\\,100546. While signs of ongoing accretion indicate the star's infancy \\citep{Deleuil2004}, the estimated age of 10\\,Myr makes it unusually old for a star with a disk \\citep{vdAncker1997}. HD\\,100546 is nearby (103 pc) and optical/near-IR scattered light imaging has revealed a wealth of structures in the disk \\citep[e.g.,][]{Grady2001, Augereau2001, Ardila2007}. Based on the  mid-- to near-IR excess ratio, \\citet{Bouwman2003} suggested an inner cavity in the disk. Later observations confirmed this gap \\citep[e.g.,][]{Grady2005, Benisty2010}. This remarkable feature and the longevity of the HD\\,100546 disk may point to a young planet in the inner 10 AU, making it a prime target for detailed studies.  The system of HD\\,100546 has a rich 2.4 -- 180~$\\mu$m spectrum as observed with ISO \\citep{Malfait1998}, showing a striking similarity with that of the comet Hale-Bopp \\citep{Crovisier1997}. \\citet{Juhasz2010} investigated the \\emph{Spitzer} spectrum of the source. Strong forsterite emission is observed in both ISO and \\emph{Spitzer} data. A continuum fit indicated crystalline dust at 210\\,K and 50\\,K \\citep{Malfait1998}.  The ISO spectrum shows a strong line of [O~{\\sc i}] 63\\,$\\mu$m and a weaker [C~{\\sc ii}] 158\\,$\\mu$m line. In the near-IR, \\citet{Brittain2009} and \\citet{vdPlas2009} observed CO ro--vibrational lines with a rotational temperature of $\\sim$1000\\,K and found no CO gas in the inner dust cavity. Observations of optical [O {\\sc i}] and Balmer line emission, on the other hand, demonstrate that the inner disk is not completely devoid of gas \\citep{AckevdAncker2006}. At submillimetre wavelengths, \\citet{Panic2010} have detected pure rotational lines of CO up to $J=7-6$ probing a warm (60--70\\,K) layer in the outer disk ($\\sim$\\,100\\,AU). Information on the intermediate temperatures is still lacking.   Here we present an analysis of narrow features (about 1\\,$\\mu$m) in the far-IR spectrum of HD\\,100546 over the full PACS range that was obtained within the \\emph{Herschel} key program `Dust, Ice and Gas in Time' (DIGIT). ",
        "conclusions": "\\label{sec:discussion} The PACS data provide the strongest constraints yet on the composition of olivines in a protoplanetary disk. We found the olivines to be extremely iron-poor (less than 3--4\\% iron). Our observations can be modeled with dust at $\\sim$70\\,K, consistent with ISO results, but requiring at least 2\\% iron. On the other hand the determination of the temperature of the forsterite based on intrinsic features of the 69\\,$\\mu$m band instead of using a continuum analysis offers a second option: The emission may emerge from pure forsterite at 200\\,K and 13\\,AU close to the midplane. This region is optically thick for shorter wavelengths.  The far-infrared spectrum of HD\\,100546 contains a wealth of molecular gas lines including CO, H$_2$O, and OH. We found CO transitions up to $J$=31--30 which arise from gas in the temperature range of 300 to 800 K and appear to sample gas in the surface layers of the disk at temperatures and surface densities intermediate to those probed by submillimeter and near-infrared CO measurements."
    },
    "1005/1005.2762.txt": {
        "abstract": "\\noindent It has been recognized that the turbulent cross helicity (correlation between the velocity and magnetic-field fluctuations) can play an important role in several magnetohydrodynamic (MHD) plasma phenomena such as the global magnetic-field generation, turbulence suppression, etc. Despite its relevance to the cross-helicity evolution, little attention has been paid to the dissipation rate of the turbulent cross helicity, $\\varepsilon_W$. In this paper, we consider the model expression for the dissipation rate of the turbulent cross helicity. In addition to the algebraic model, an evolution equation of $\\varepsilon_W$ is proposed on the basis of the statistical analytical theory of inhomogeneous turbulence. A turbulence model with the modeling of $\\varepsilon_W$ is applied to the solar-wind turbulence. Numerical results on the large-scale evolution of the cross helicity is compared with the satellite observations. It is shown that, as far as the solar-wind application is concerned, the simplest possible algebraic model for $\\varepsilon_W$ is sufficient for elucidating the large-scale spatial evolution of the solar-wind turbulence. Dependence of the cross-helicity evolution on the large-scale velocity structures such as velocity shear and flow expansion is also discussed.  \\vspace{6pt}  \\noindent {{\\bf{Keywords:}} Magnetohydrodynamic turbulence; turbulence model; cross helicity; dissipation rate; solar wind; }\\bigskip ",
        "introduction": "\\label{sec:level1} In the magnetohydrodynamic (MHD) turbulent flow at high magnetic Reynolds number ($Rm \\gg 1$), magnetic fields are considered to be frozen in plasmas, and move with the flow.\\cite{alf1950} In such a flow, the induced magnetic field is often much larger than the originally imposed field. Besides, MHD waves such as the Alfv\\'{e}n wave are considered to exist ubiquitously. The cross helicity, defined by the correlation between the velocity $\\bf{u}$ and magnetic field $\\bf{b}$, is a possible describer of such MHD turbulence properties. Actually, the magnetic-field generation due to the turbulent cross helicity has been investigated.\\cite{yos1990,yos1993,yok1996,yos1998,yos1999,yok1999,yos2000,yos2004}  As is well known, the total amount of cross helicity $\\int_V {{\\bf{u}} \\cdot {\\bf{b}}} dV$, as well as that of the MHD energy $\\int_V {({\\bf{u}}^2 + {\\bf{b}}^2) / 2}\\ dV$, is an inviscid invariant of the MHD equations. Because of this conservative property, the turbulent densities of the MHD energy and cross-helicity, $K \\equiv \\langle { {\\bf{u}}' {}^2 + {\\bf{b}}' {}^2 } \\rangle / 2$ and $W \\equiv \\langle { {\\bf{u}}' \\cdot {\\bf{b}}' } \\rangle$, may serve themselves as a good measure for characterizing the statistical properties of MHD turbulence (${\\bf{u}}'$: velocity fluctuation, ${\\bf{b}}'$: magnetic-field fluctuation, $\\langle \\cdots \\rangle$: ensemble average).  The evolution equations of $K$ and $W$ are similar in form, and their mathematical structures are quite simple. The evolution of $K$ and $W$ are determined by three constitutes: the production, dissipation, and transport rates. Firstly, the production rate is expressed by the correlation of turbulence fields coupled with the mean-field inhomogeneity. We note that the production rate of turbulence quantities such as $\\langle {{\\bf{u}}' \\cdot {\\bf{b}}'} \\rangle$ can be expressed exactly in the same as the counterpart of mean-field quantities such as ${\\bf{U}} \\cdot {\\bf{B}}$, but with the opposite sign ($\\bf{U}$: the mean velocity, $\\bf{B}$: the mean magnetic field). This makes our interpretation possible that the drain of mean-field quantity gives rise to the generation of turbulence counterpart. So, the production rate represents how a quantity is supplied to turbulence by way of its cascading process (See Appendix~\\ref{sec:appendixA}). Secondly, the dissipation rates of the turbulent MHD energy and cross helicity, $\\varepsilon$ and $\\varepsilon_W$, whose definitions will be given shortly in Section~\\ref{sec:level2}, represent the effects of molecular viscosity and magnetic diffusivity coupled with the small-scale fluctuations. However, we stress the following point. The dissipation rates of the turbulent MHD energy and cross helicity, $\\varepsilon$ and $\\varepsilon_W$, can be considered from another aspect. In the intermediate range of turbulence, called the inertial range, the energy and cross helicity supplied from the energy-containing range compensate the energy and cross-helicity lost in the dissipation range. For this cascade picture of turbulence, the energy and cross-helicity transfer from lower to higher wavenumber ranges are most important quantities. In an equilibrium turbulence, $\\varepsilon$ and $\\varepsilon_W$ represent these transfer rates of $K$ and $W$, respectively. This makes the construction of the $\\varepsilon$ and $\\varepsilon_W$ equation possible as we show later. Finally, the transport rates express the flux of a quantity that enters the fluid volume through the boundary. The expression for the transport rates suggests in what situation the quantities considered can be supplied to turbulence.  Thanks to these clear-cut pictures associated with the evolution equation, the cross helicity (density) $W$, as well as the turbulent MHD energy (density) $K$, may play an important role in the turbulence modeling of MHD fluids. However, as compared with $K$ and other pseudoscalar turbulence quantities such as the turbulence kinetic and magnetic helicities, only a limited attention so far has been paid to the cross helicity.  In the context of homogeneous isotropic MHD turbulence, some important investigations have been made on the decaying rate of the cross helicity or $\\varepsilon_W$. It was shown that if there is a prevailed sign of the cross helicity in the initial state, the system goes towards a dynamically aligned state. The cross helicity scaled by the MHD energy grows towards +1 or -1 depending on the initially prevailed sign of the cross helicity.\\cite{dob1980a,dob1980b,gra1982,gra1983}  In the context of inhomogeneous MHD turbulence, the cross helicity has been investigated mostly in the solar-wind research. By using spacecraft observations, detailed spectra of cross helicity have been examined.\\cite{bel1971,rob1987,tum1995} In order to explain the large-scale behavior of the solar-wind turbulence, several models have been proposed.\\cite{zho1990,tum1995} However, investigations related to the cross helicity are mostly concentrated on arguments of its production rate, and effects of large-scale inhomogeneities such as the mean velocity shear have been discussed. Matthaeus and coworkers have employed a kind of algebraic model for the cross-helicity dissipation rate.\\cite{zho1990,mat2004,bre2005,bre2008} Adopting this algebraic model of $\\varepsilon_W$, Usmanov {\\it et al}.\\ have recently performed a series of elaborated numerical simulations on the large-scale evolution of solar-wind turbulence.\\cite{usm2011} However, generally speaking, arguments concerning the dissipation rate of $W$ are still far from sufficient.  Also in the context of the turbulence dynamo, the transport equation for the turbulent cross helicity has been considered, where an algebraic model for the cross-helicity dissipation rate has been proposed.\\cite{yos1990,yos1993} Sur and Brandenburg wrote down an evolution equation for the cross-helicity effect, and argued the cross-helicity destruction in the context of quenching mechanism.\\cite{sur2009}  In order to examine the evolution of the turbulent cross helicity $W$, it is indispensable to properly estimate the dissipation rate of $W$, $\\varepsilon_W$, as well as the cross-helicity production rate $P_W$. We address this problem; modeling the cross-helicity dissipation rate on the basis of a statistical analytical theory.  Spacecraft observations of solar-wind turbulence have revealed detailed information on the large-scale behavior of turbulent statistical quantities, which includes the radial evolution of the cross helicity both in the low- and high-speed wind regions. Comparison of the satellite observations with the numerical simulation with the aid of a turbulence model provides a good test for the cross-helicity dissipation models. In this work, we will delve into the problem of cross-helicity dissipation modeling by using such comparisons.  The organization of this paper is as follows. After briefly showing the evolution equation of the cross helicity in Section~\\ref{sec:level2}, we present the exact equation of the cross-helicity dissipation rate $\\varepsilon_W$ in Section~\\ref{sec:level3}. Then, we consider two candidates for the model of $\\varepsilon_W$ in Section~\\ref{sec:level4}; one is the algebraic model and another is a transport-equation model. These expressions are systematically derived with the aid of a statistical analytical theory of inhomogeneous turbulence. Some features of the adopted turbulence model, which is an expansion of the hydrodynamic $k - \\epsilon$-type one-point turbulence model in the engineering field, are noted in Section~\\ref{sec:level5}. An application of the model to the solar wind is presented in Section~\\ref{sec:level6}. A brief summary is given in Section~\\ref{sec:level7}.    %------------------------------------------------------------- %\t2 Equation for the cross helicity %------------------------------------------------------------- ",
        "conclusions": "\\label{sec:level7} The evolution of turbulent cross helicity ($W \\equiv \\langle {{\\bf{u}}' \\cdot {\\bf{b}}'} \\rangle$) was investigated from the viewpoint of generation and destruction mechanism of it. In particular, two possibilities of expressing the dissipation rate of $W$, $\\varepsilon_W$, were presented; (i) the algebraic model and (ii) the evolution equation for $\\varepsilon_W$. It was shown that both model expressions can be systematically derived from the zeroth- and first-order calculations of the statistical analytical theory of inhomogeneous turbulence, respectively. Validity of the model expressions was examined with the aid of a turbulence model constituted by four one-point turbulent statistical quantities (the turbulent MHD energy $K$, its dissipation rate $\\varepsilon$, the turbulent cross helicity $W$, and the turbulent residual energy $K_{\\rm{R}}$). It was shown that, as far as the application to the solar-wind turbulence is concerned, the algebraic model for $\\varepsilon_W$ with one model constant gives results plausible enough. Dependence of the cross-helicity evolution on the large-scale velocity structure was also discussed. In the context of the solar wind, it was shown that both of flow expansion and large-scale velocity shear contribute to decreasing the magnitude of the scaled cross helicity, which confirmed earlier results.\\cite{mat2004,bre2005,bre2008} More detailed expressions for the $\\varepsilon_W$ equation were also indicated from the higher-order calculation of the statistical analytical theory of inhomogeneous turbulence.   %------------------------------------------------------------- %\t\tAcknowledgments %-------------------------------------------------------------"
    },
    "1005/1005.2538_arXiv.txt": {
        "abstract": "We determine an expression for the cosmic variance of any ``normal'' galaxy survey based on examination of $M^* \\pm 1$ mag galaxies in the SDSS DR7 data cube. We find that cosmic variance will depend on a number of factors principally: total survey volume, survey aspect ratio, and whether the area surveyed is contiguous or comprised of independent sight-lines. As a rule of thumb cosmic variance falls below 10\\% once a volume of $10^7h_{0.7}^{-3}$Mpc$^3$ is surveyed for a single contiguous region with a 1:1 aspect ratio. Cosmic variance will be lower for higher aspect ratios and/or non-contiguous surveys. Extrapolating outside our test region we infer that cosmic variance in the entire SDSS DR7 main survey region is $\\sim 7$\\% to $z < 0.1$  The equation obtained from the SDSS DR7 region can be generalised to estimate the cosmic variance for any density measurement determined from normal galaxies (e.g., luminosity densities, stellar mass densities and cosmic star-formation rates) within the volume range $10^3$ to $10^7 h^{-3}_{0.7}$Mpc$^3$.  We apply our equation to show that 2 sightlines are required to ensure cosmic variance is $<10$\\% in any ASKAP galaxy survey (divided into $\\Delta z \\sim 0.1$ intervals, i.e., $\\sim 1$ Gyr intervals for $z <0.5$). Likewise 10 MeerKAT sightlines will be required to meet the same conditions. GAMA, VVDS, and zCOSMOS all suffer less than 10\\% cosmic variance ($\\sim$ 3\\%-8\\%) in $\\Delta z$ intervals of 0.1, 0.25, and 0.5 respectively. Finally we show that cosmic variance is potentially at the 50-70\\% level, or greater, in the HST Ultra Deep Field depending on assumptions as to the evolution of clustering. 100 or 10 independent sightlines will be required to reduce cosmic variance to a manageable level ($<10$\\%) for HST ACS or HST WFC3 surveys respectively (in $\\Delta z \\sim 1$ intervals). Cosmic variance is therefore a significant factor in the $z>6$ HST studies currently underway. ",
        "introduction": "The Universe is not homogeneous except on the largest scales ($>1$Gpc, Davis et al.~1985). As a consequence number and density measurements derived from within modest volumes will show greater than Poisson variation (see Szapudi \\& Colombi 1996 for example). This cosmic variance\\footnote{Technically the term sample variance is more correct but here we adhere to the current convention of using the term cosmic variance to describe perturbations in measurements within our Universe due to sampling size.}, or small-scale scale-dependent inhomogeneity, is often the dominant source of error in many contemporary extragalactic measurements. Examples include the galaxy luminosity function/densities (Norberg et al.~2002a; Hill et al.~2010), the HI mass function (Zwaan et al.~2005), the cosmic star-formation history (Hopkins \\& Beacom~2006), and the stellar mass density (Wilkins, Trentham \\& Hopkins~2008). Generally any number or density measurement derived from the galaxy population as a whole is susceptible. Typically, although often neglected in many studies, the cosmic variance can be estimated through one of four methods: Comparison with numerical simulations which encompass larger volumes such as the Millennium Simulation (e.g., Newman \\& Davis~2002; Somerville et al.~2004; Trenti \\& Stiavelli~2008; Moster et al.~2010); analytically using measurements of the 2 or 3-pt correlation functions (e.g., Driver et al.~2003); empirically by Monte-Carlo sampling of a larger survey (e.g., Driver et al.~2005; Hill et al.~2010); or, also empirically, by Jackknife sampling of the volume in question (e.g., Liske et al.~2003). These methods all have strengths and weaknesses can be laborious to impliment and potentially inconsistent depending on the method adopted and the assumptions made. For example to estimate the cosmic variance from numerical simulations requires the adoption of a numerical simulation (e.g., Springer et al.~2005), and either a semi-analytical prescription (Cole et al.~ 2000; Baugh~2006), or a halo occupation distribution (Berlind \\& Weinberg~2002; Moster et al.~2010) before an appropriate cosmic variance estimate can be made (e.g., Moster et al.~2010). The analytical method (e.g., Driver et al.~2003) implicitly assumes Poisson statistics, radial symmetry, and (because it is parametric) smoothes over potential 'features' in the underlying distribution (e.g., BAOs). The empirical method is not always practical if a large suitable survey does not exist, and Jackknife sampling (where one divides the sample into many parts and recomputes the value in question with each part missing in turn) is only capable of revealing the cosmic variance on scales smaller than the volume in question (nevertheless a useful indicator as cosmic variance should generally decrease with increasing scale).  In this paper we aim to use the largest volume survey to date, the Sloan Digital Sky Survey (SDSS), to empirically determine some generically useful formulae for estimating the cosmic variance as a function of survey volume and survey shape. These formulae should also assist in the design of future surveys where trade offs between area and depth need to be made. Throughout we adopt a standard cosmology with the following parameter set: $\\Omega_M=0.3, \\Omega_{\\Lambda}=0.7, H_o=70$kms$^{-1}$Mpc$^{-1}$ although as our analysis is based on very local volumes only the value of the adopted Hubble constant is significant. ",
        "conclusions": "We have derived a simple empirical expression for calculating cosmic variance for almost any extragalactic survey. The results are entirely empirical and based on resampling the SDSS DR7. The resulting equations agree extremely well with the recent numerical results by Moster et al.~(2010). The two resulting equations provide corrections for $z<0.1$ robustly and for $z>0.1$ under the following caveats:  ~  \\noindent (1) Te derived cosmic variance is for $M^* \\pm 1$ mag population only and assumed not to evolve with lookback time  -- this is clearly incompatible with our understanding of the evolution of structure and hence beyond $z \\sim 1$ the derived values should be taken as indicative only.  ~  \\noindent (2) That above $250h^{-1}_{0.7}$Mpc cosmic variance scales with radial co-moving length according to Poisson statistics.  ~  The two equations are then used to determine cosmic variance values for a number of recent, ongoing and planned surveys."
    },
    "1005/1005.0684_arXiv.txt": {
        "abstract": "We analysed the IGR~J16465--4507 Burst Alert Teelescope survey data collected during the first 54 months of the \\sw\\ mission. The source is in a crowded field and it is revealed through an ad hoc imaging analysis at a significance level of $\\sim$14 standard deviations. The 15--50 keV average flux is $\\sim3 \\times 10^{-11}$ \\ferg. The timing analysis  reveals an orbital period of 30.243$\\pm$0.035 days. The folded light curve shows the presence of a wide phase interval of minimum intensity, lasting $\\sim 20\\%$ of the orbital period. This could be explained with a full eclipse of the compact object in an extremely eccentric orbit or with the passage of the compact source through a lower density wind at the orbit apastron. The modest dynamical range  observed during the BAT monitoring suggests that IGR~J16465$-$4507 is a wind-fed system, continuously accreting from a rather homogeneous wind, and not a member of the Supergiant Fast X-ray Transient class. ",
        "introduction": "}   \\begin{figure}% \\begin{center} \\centerline{\\includegraphics[width=7.5cm,angle=0]{figure1a.ps}} \\centerline{\\includegraphics[width=6cm,angle=270]{figure1b.ps}} \\caption[IGRJ16465-450 sky map]{{\\bf Top}: 15--50 keV significance map in the neighborhood of IGR~J16465$-$4507. The color-bar represents the significance levels. The cyan star marks the XMM position \\citep{ZuritaHeras2004:16465-4507} of IGR J16465$-$4507. The white circles are centered on the position derived by fitting the significance profile extracted along the white line with two Gaussians plus a constant; their radius corresponds to the 90\\% error on the position. {\\bf Bottom}: Significance profile extracted along the white line in the above map. The plot shows the data (stars) and the best fit model (blue line, the sum of two Gaussian profiles plus a constant value). The higher peak (red line) corresponds to IGR J16479$-$4514, the lower peak (green line) corresponds to IGR J16465$-$4507. } \\label{map} \\end{center} \\end{figure}   High mass X--ray binaries (HMXBs), stellar systems composed of a compact object and an early-type massive star, are traditionally divided in two subclasses \\citep[e.g.][and references therein]{vanParadijs1995:binaries}, depending on the nature of the high mass primary and, consequently, the different mass-transfer and accretion mechanism. On one side are the systems with main sequence Be primaries (Be-HMXBs). They are generally wide ($P_{\\rm orb}\\ga 10$\\,d) eccentric (eccentricity $e\\sim 0.3$--0.5) systems in which the primaries are not filling their Roche lobe, and accretion onto the compact object occurs from the equatorial region of the rapidly rotating Be star. Most of these systems are highly variable: in some of them recurrent outbursts are observed caused by an enhanced rate when the compact star passes close to the Be star. On the other side are the systems with an evolved OB supergiant primary (sgHMXB). Their periods are shorter ($P_{\\rm orb}\\la 10$\\,d) and their orbits more circular than in Be-HMXBs. They are powered either by a geometrically thin accretion disc or by the strong radiation-driven stellar wind, depending on whether the primary fills its Roche lobe or not. Their X-ray emission is bright and persistent.  Recently, this rather clear-cut picture was made more structured with the INTEGRAL observations of the Galactic plane. Two additional classes  were added to the classical OB primary HMXBs: the highly absorbed persistent systems (Walter et al. 2004, 2006) and the supergiant fast X--ray transients \\citep[SFXTs, ][]{Sguera2005,negueruela06esa, smith04,intzand05}. The former are characterized by orbital and spin periods consistent with those observed in wind-accreting systems, but a much higher absorbing column density. The latter are transient sources showing a large dynamic range of 3--5 orders of magnitude with sporadic outbursts (which however are significantly shorter than those of typical Be-HMXBs), characterized by bright flares lasting up to days with peak luminosities of 10$^{36}$--10$^{37}$~erg~s$^{-1}$ \\citep{Sguera2005,Romano2009:sfxts_paperV,sidoli09}.  The Burst Alert Telescope \\citep[BAT,][]{Barthelmy2005:BAT} on board \\sw\\ \\citep{Gehrels2004mn} is performing a continuous coverage of the hard X-ray sky (50 to 80\\% of the sky every day). This allowed the detection of many of the new INTEGRAL HMXBs (e.g. \\citealp{cusumano10}) and the collection of their long term light curves. In this Letter we analyse the hard X-ray data collected  during the first 54 months of \\sw-BAT sky monitoring using data in the region of the IGR~J16465--4507. This source was discovered by INTEGRAL in 2004 \\citep{Lutovinov2004:16465-4507}  and X-ray activity was observed with IBIS/ISGRI starting on September 6, at a flux level of 8.8 $\\pm$ 0.9 mCrab (18-60 keV), followed by a flare (up to 28 mCrab) on September 7. Follow-up observations with XMM-Newton revealed pulsations at 228$\\pm$6 s \\citep{Lutovinov2005}  and allowed the identification of the optical counterpart with 2MASS~J16463526--4507045 \\citep{ZuritaHeras2004:16465-4507}. This was classified as a B0.5 Ib supergiant at a distance of $\\sim 8$\\,kpc \\citep{Negueruela2007} or as a O9.5 Ia supergiant at a distance of $9.5_{-5.7}^{+14.1}$\\,kpc \\citep{Nespoli2008}. The supergiant nature of the companion, combined with the observed hard X-ray variability \\citep{Lutovinov2004:16465-4507} and the X-ray spectral distribution, modeled by a hard power law with photon index 1.0$\\pm$0.52 \\citep{Lutovinov2005}, suggested a classification of this source as a SFXT \\citep{Negueruela2006}. However, \\citet{Walter2007}, based on  INTEGRAL measurements, suggested that  IGR~J16465--4507 is likely a classical supergiant HMXB, with an average flux just below the IBIS/ISGRI sensitivity undergoing sporadic long periods of enhanced activity.  This Letter is organized as follows. Section 2 describes the BAT data reduction and the imaging analysis. Section 3 reports on the timing analysis. Sect. 4 describes the analysis of the pointed soft X-ray observation with \\sw-XRT. In Sect.\\ 5 we briefly discuss our results. Errors are at 90\\,\\% confidence level, if not stated otherwise. ",
        "conclusions": "}  We have analysed the data collected by \\sw-BAT during the first 54 months of the \\sw\\ mission in the region of the supergiant HMXB \\src16465 {}. The source is detected at a significance level of $\\sim$ 14 standard deviations with an average flux of $\\sim 2.36 \\times 10^{-11}$ \\ferg\\ in the 15--50 keV energy band.  The light curve reveals a periodicity of $30.243 \\pm 0.035$ days that we interpret as  the orbital period of the binary system.  By applying Kepler's third law, the semi-major axis of the binary system is given by $a^3=P_{\\rm orb}^2\\times \\rm G(M_{\\star}+M_{\\rm X})/4\\pi^2$, where M$_{\\star}$ and M$_{\\rm X}$ are the masses of the supergiant and compact object, respectively. We adopt M$_{\\rm X}=1.4$~M$_{\\sun}$ and M$_{\\star}=27.8$~M$_{\\sun}$ \\citep[][for an O9.5 I star of radius R$_{\\star}=22.1$~R$_{\\sun}$]{Martins2005}. This yields $a\\sim125$~R$_{\\sun} \\sim 6$~R$_{\\star}$. We note that the assumption of a stellar type B0.5 Ib would lead to an estimate of $a\\sim150$~R$_{\\sun} \\sim 5$~R$_{\\star}$ \\citep[][M$_{\\star}=47$~M$_{\\sun}$, R$_{\\star}=32.2$~R$_{\\sun}$]{Searle2008}.  The folded profile shows the presence of a dip with a count rate consistent with no emission, that could be intepreted as a full eclipse. However, the width of this dip ($\\sim 20\\%$ of P$_{\\rm orb}$) is not consistent with the duration of the eclipse expected for a circular orbit with 6 R$_{\\star}$ semimajor axis, assuming an edge-on inclination. To obtain such a long eclipse, the system should have an eccentricity of at least 0.8 coupled with a high inclination. This scenario is unlikely. Alternatively, this wide phase interval of minimum intensity could be explained with an eccentric orbit and a lower wind density at apastron.    Given our knowledge of both a spin period and an orbital period we can locate the source on the Corbet diagram \\citep{Corbet1986:diagram} together with the known $P_{\\rm spin}$ and $P_{\\rm orb}$ of other binary systems \\citep{Bildsten1997,Liu2006:hmxb}. We also plot in the diagram the position of the two SFXTs IGR J18483-0311 \\citep{Levine2006:igr18483,Zurita2009:sax1818.6_period} and IGR~J11215-5952 \\citep{Swank2007:atel999,Sidoli2007,Romano2009:11215_2008}. \\src16465 {} sits at the boundary of the wind-fed OB-HMXBs (systems with OB supergiants that underfill their Roche-lobes) and the locus of the Be transients (Fig.~\\ref{corbet}.  The characteristics of the system \\citep{Lutovinov2004:16465-4507,Lutovinov2005} suggested a classification of this source as a SFXT \\citep{Negueruela2006}. However, the outburst history of this source is rather scarce. After the initial discovery by \\inte\\ no further outbursts have been reported. \\citet{Walter2007} mention three episodes of enhanced hard X-ray activity observed by \\inte\\ at the limit of the instrumental sensitivity, that cannot be considered as real flares. The BAT 15--50\\,keV light curve during the 54 months of monitoring indicates that this is a  faint persistent source, with an average luminosity of $2.6\\times10^{35}$ erg s$^{-1}$ (at a distance of 9.5 kpc), no evidence for flaring activity and a weak variability with a dynamical range lower than 10.   Based on their X-ray variability, sgHMXB can be classified as classical and absorbed systems ( variability factor $<$20) or as SFXT (variability factor $>100$). The timing behavior of IGR~J16465$-$4507 suggests a wind-fed system, with a neutron star continuously accreting from a rather homogeneous wind. The luminosity, lower than what observed in classical sgHMXB ($10^{36}-10^{37}$ erg s$^{-1}$), can be explained with the larger orbital separation ($\\sim$5 R$_{\\star}$ compared to $\\sim$2 R$_{\\star}$ in classical systems).     \\begin{figure}% \\begin{center} \\centerline{\\includegraphics[width=6cm,angle=90]{figure5.ps}} \\caption[]{The Corbet diagram showing the neutron star $P_{\\rm spin}$ vs.\\ binary period $P_{\\rm orb}$. Black circles are HMXBs with an OB primary, empty circles those with a Be one. Larger (red) symbols represent SFXTs: triangles are IGR~J18483$-$0311  and IGR~J11215$-$5952, the star the newly determined position of \\src16465 . } \\label{corbet} \\end{center} \\end{figure}"
    },
    "1005/1005.2681_arXiv.txt": {
        "abstract": "{Neutrinos emitted from a supernova encode useful information about neutrino physics and astrophysics. Interpreting the neutrino signal depends crucially on understanding neutrino production, flavor mixing during propagation, and detection. In this talk, we review the physics potential of a SN neutrino observation.}  \\FullConference{35th International Conference of High Energy Physics - ICHEP2010,\\\\ July 22-28, 2010\\\\ Paris France}   \\begin{document} ",
        "introduction": "Neutrinos from a core-collapse supernova (SN)~\\footnote{Neutrinos refer to both neutrinos and antineutrinos. Only neutrinos will be written as $\\nu$. Similarly, antineutrinos will be written as $\\bar{\\nu}$. Also, SN refers to core-collapse supernovae only.} provide a rare and valuable physics opportunity~\\cite{Dighe:2008dq}. The gargantuan fluxes and low interaction rates of SN neutrinos, makes them ideal candidates to probe neutrino mixing properties and study the extreme conditions in the depths of a star.  Observation of the SN $1987$A in neutrinos opened the field of neutrino astronomy~\\cite{Koshiba:1992yb}, and confirmed our overall understanding of SN neutrino physics~\\cite{Raffelt:1996wa}. With present and planned detectors~\\cite{Scholberg:2010zz}, a galactic SN is expected to result in a high-statistics detection. This would lead to  significant advances in neutrino physics and SN astrophysics, if we can disentangle the relevant information.  Decoding the neutrino signal requires a detailed understanding of neutrino production, mixing during propagation, and detection. For neutrinos, the mixing scenario is reasonably well-determined, except the value of $\\theta_{13}$, the sign of $\\Delta m^2_{\\rm atm}$, and the CP-violating phase~\\cite{GonzalezGarcia:2010er}, and the main detection channels are well-calibrated. However for SNe, initial neutrino fluxes and spectra predicted from SN theory have a significant variance~\\cite{Simulations}. Similarly, the stellar conditions are poorly constrained. The signal interpretation is plagued by these uncertainties, and the strategy for disentangling information from a SN neutrino signal must rely on generic features that are insensitive to model assumptions.  Some information is obtained directly, e.g. the direction and time structure of the event, and the above mentioned uncertainties do not affect our inferences.  But a lot more information, e.g. flavor dependent energy spectra of the neutrinos, clues to the unknown neutrino parameters, and some signatures of stellar dynamics, is encoded in a flavor dependent way. A detailed treatment of neutrino mixing is required to extract that information. Traditional analyses of neutrino mixing for SN took into account interactions with matter, through the Mikheev-Smirnov-Wolfenstein  (MSW) effect. This MSW-based paradigm was believed to be complete, and numerous results followed. See the review~\\cite{Dighe:2004xy} for the traditional expectations. However, that picture was incomplete. Deep in the SN, neutrino densities are large enough to make their collective interactions extremely important. The mixing angles are highly matter suppressed, and one may expect no flavor conversion in that region. However, this naive expectation is incorrect. Neutrino-neutrino interactions entangle the flavor evolution of all neutrinos and create an instability in the flavor composition. Thus flavor conversions take place even for extremely small mixing angles, with a rich phenomenology.  See the recent review~\\cite{Duan:2010bg} on these so-called  ``Collective effects''. In this talk, we summarize the main aspects of SN neutrinos, focussing on the collective effects and their impact on the physics potential of SN neutrino observations. ",
        "conclusions": "The main hurdle in interpreting SN neutrino data, besides the absence of it at present, is the lack of knowledge of initial conditions, i.e. the initial fluxes and densities. This is likely to lead to degeneracies. However, one may expect various aspects of the time and flavor dependent SN signal to be used in synergy. Future developments in SN theory/simulations could be expected to reduce or eliminate some of these degeneracies. Another area of improvement would be a better understanding of the flavor conversion. In particular, the effect of anisotropy and inhomogeneity on collective effects and MSW conversions.  A future galactic SN can be expected to provide a wealth of scientific information for neutrino oscillation physics and SN astrophysics. This is a rare opportunity, occurring once-in-a-lifetime, and we must be ready with suitable detectors and the required theoretical understanding to interpret the data. A significant step towards this, would be a detailed understanding of the rich phenomenology of neutrinos from supernovae."
    },
    "1005/1005.5068_arXiv.txt": {
        "abstract": "Radio pulsar surveys are producing many more pulsar candidates than can be inspected by human experts in a practical length of time.  Here we present a technique to automatically identify credible pulsar candidates from pulsar surveys using an artificial neural network. The technique has been applied to candidates from a recent re-analysis of the Parkes multi-beam pulsar survey resulting in the discovery of a previously unidentified pulsar. ",
        "introduction": "Since the discovery of pulsars by Jocelyn Bell and Antony Hewish at Cambridge in 1967 using a pen chart recorder (Hewish et al. 1968\\nocite{hbp+68}), pulsar searching has come a long way. Modern pulsar surveys use high performance computing facilities to perform an extensive range of signal processing and search algorithms. These methods are designed to maximize sensitivity to weak, rapid, and dispersed pulsar signals often buried in large amounts of terrestrial radio frequency interference (RFI), or even in binary systems. There is little doubt that these complex algorithms have aided searches for pulsars, however there remain certain search tasks for which standard computer programs are of little use. In particular, the final stage of a pulsar search, the selection of credible pulsar candidates for follow-up observations, which still remains a task for a human since the decision is visual and based on a number of combined properties of the pulsar signal. The process can be time consuming and inefficient in analysis of large-scale surveys that produce many millions of pulsar candidates.  Large-scale pulsar surveys such as the Parkes multi-beam pulsar survey (PMPS) (Manchester et al. 2001)\\nocite{mlc+01} have dramatically increased the number of known pulsars. Finding more pulsars elucidates the properties of their Galactic population, and also offers the possibility of uncovering new and extreme phenomena in neutron-star astrophysics. Future pulsar surveys will be done with the next generation of radio telescopes, such as LOw Frequency ARray (LOFAR), the Five hundred metre Aperture Spherical Telescope (FAST), and the Square Kilometre Array (SKA) (e.g. van Leeuwen \\& Stappers 2010\\nocite{2010A&A...509A...7V}, Smits et al., 2009a,b\\nocite{2009A&A...505..919S}\\nocite{2009A&A...493.1161S}, Cordes et al., 2004\\nocite{2004NewAR..48.1413C}).  These radio telescopes will be excellent survey tools because of their large collecting areas, capability to form many simultaneous beams on the sky, and in the case of the interferometers, wide fields of view.  It is expected that these instruments will detect a large fraction of the observable pulsars in the Galaxy. The inevitable flood of pulsar candidates that will require inspection to achieve this will certainly require some form of multi-person or machine based candidate selection. Some large-scale astronomical surveys and data mining projects have resorted to employing many online volunteers to search for or classify their objects of interest in so-called `citizen science' projects (e.g. Lintott et al., 2008\\nocite{2008MNRAS.389.1179L}, Westphal et al., 2005\\nocite{2005LPI....36.1908W}). In searches for pulsars, the pulsar Arecibo L-band Feed Array (ALFA) survey collaboration has enlisted the help of High School and undergraduate students in a successful outreach and science program to identify potential pulsar candidates (Jenet et al., 2007\\nocite{2007AAS...211.0517J}).  Recent machine solutions include, candidate ranking based on likelihoods calculated from parameter distributions of pulsar and non-pulsar signals (Lee, private communication), and the sorting of candidates based on a number of `scores' that indicate the similarity of the signal to that of a typical pulsar (Keith et al., 2009)\\nocite{kel+09}. In this paper we present an alternative method whereby an artificial neural network (ANN) has been trained using a particular set of scores to automatically identify credible pulsar candidates from a recent re-analysis of the PMPS. ANNs have long been used in other areas of astronomy, for example in the morphological classification of galaxies (e.g. Storrie-Lombardi et al., 1992\\nocite{1992MNRAS.259P...8S}, Zhang, Li \\& Zhao 2009\\nocite{2009MNRAS.392..233Z}), the estimation of photometric redshifts of sources in the Sloan Digital Sky Survey (Firth, Lahav \\& Somerville 2003\\nocite{2003MNRAS.339.1195F}) and in the selection of microlensing events from large variability surveys \\nocite{2003MNRAS.341.1373B} (Belokurov, Evans \\& Du 2003).  The outline of this paper is as follows: In Section 2 we describe how pulsar candidate selection in our recent re-analysis of the PMPS is typically done and some of the problems associated with this method. Section 3 gives a brief introduction to ANNs. In Section 4 we describe our implementation of an ANN to classify pulsar candidates from our re-analysis of the PMPS and the results from a test on a sample of the search output data.  Preliminary parameters of the pulsar discovered using an ANN during our re-analysis of the PMPS are given in Section 5. Finally Section 6 gives a summary, followed by a discussion of the future application of ANNs to pulsar surveys. ",
        "conclusions": "ANNs have been developed to aid searches for radio pulsars in a database of $\\sim$ 16 million pulsar candidates generated from a recent re-analysis of the PMPS. The ANNs were trained using small sets of characteristic scores derived from the pulsar candidate plots generated for human inspection. Around $92$ per cent of the pulsars present in a test sample of $\\sim$ 2.5 million pulsar candidates can be recovered by our ANNs.  It is likely that pulsars from the test sample were missed due to one of three reasons: poor training of the ANNs on MSPs, abnormal candidate plots generated by our search software, or unbalanced training sets.  Future implementations should avoid unbalanced training sets or bias the training sets toward the objects of interest. A larger number of false positives from the ANN is beneficial for searches of the inherently rare objects in the data set viz. pulsars \\nocite{2003MNRAS.341.1373B} (e.g. Belokurov, Evans \\& Du 2003).  To test if the `contamination' from non-pulsars in the results set could be further reduced the threshold for pulsar identification, $z_{1}$, was increased to $0.7$, $0.8$, and $0.9$. In each case the number of genuine pulsar to non-pulsar candidates recovered by the ANN were 1:24, 1:23, and 1:20 respectively, however the corresponding recovered fraction of pulsars were 89, 86, and 81 per cent. In a search for pulsars these few per cent could be potential new discoveries. As such higher thresholds should only be used when the number of candidates produced by the ANN are unmanageable. Improving the detection efficiency of our ANNs will be the subject of future work. Improvements in this area are likely to depend upon better representation of the two-dimensional diagnostic plots in our input vectors. These planes are particularly useful for discriminating against narrow-band and impulsive RFI.  ANNs are only capable of classifying the inputs given to them based on the training received. By training the ANN using a set of pulsars with particular characteristics we reduce the possibility of uncovering atypical, unusual, or unexpected phenomena. Therefore, it is important to use training sets of pulsars with a wide variety of properties and displaying phenomena such as scintillation, scattering, intermittency and binary motion. Without training on pulsars that exhibit such properties an ANN will only select against these objects. As described in Section~\\ref{s:tests} MSPs display somewhat different candidate plots to their standard pulsar counterparts. Separate ANN training using only data from MSPs may be required to find these pulsars. For more exotic pulsars, such as relativistic binary pulsars, any training sample will be limited in size. To find these systems it might be possible to form training sets from simulated pulsar signals covering a wide range of pulsar and binary parameters. However, it is expected that ANNs trained with simulated training sets will not perform as well as those based on real data due to the subtle effects of instrumentation and RFI environments. In new pulsar surveys where the search data may vary from that of previous surveys, for example in the number of frequency channels and the sampling time, dedicated observations of known pulsars will be required to build a training database of MSPs and the standard population of pulsars.  A complementary class of ANNs are unsupervised ANNs which require no desired or `target' vector during training. The best example of such a ANN is the Kohonen Self Organizing Map (SOM) (e.g. Kohonen, 2001)\\nocite{Kohonen01}. SOMs provide a way of representing multi-dimensional data in a lower number of dimensions, generally the two-dimensions of a map. Similar input vectors share similar regions on these maps. Such a ANN could be used to divide the pulsar candidates into sub-classes such as MSPs, standard pulsars, noise signals, and even common RFI signals. Although not optimal for candidate selection such ANNs will be investigated in future work.  Future pulsar surveys will probe larger volumes of both real space and the parameter space associated with radio pulsars. The SKA will have a sensitivity that will allow the discovery of all $20\\;000$ to $30\\;000$ pulsars beaming towards Earth and visible from the short listed sites in South Africa and Western Australia. Inspecting the candidates from such surveys will be a tremendous data analysis task. For illustration, scaling the PMPS up to an all sky survey would result in approximately a factor of 40 increase in the number of pulsar candidates. Using the same search algorithm as in our recent re-analysis of the PMPS and assuming an inspection time of one second per pulsar candidate, viewing all the candidates from such as survey would take over 20 years. This figure does not take into account beams that contribute many erroneous candidates due to RFI and which are relatively easy to flag and avoid using graphical selection tools. However, as we have already discussed these cuts can increase the risk of missing genuine pulsars that lie near to RFI in phase-space. Using an ANN and assuming the same recovered fraction of candidates of $0.5$ per cent the total inspection time would come down to just over one month, vastly reducing the work load on human observers.  We would like to stress that simple ANNs, like those presented in this work, should not yet be treated as a complete replacement for human inspection of pulsar candidates. Humans are still best placed to spot anything unusual or interesting about the individual pulsar candidates being viewed.  Because of the significant time savings, of around two orders of magnitude, ANNs might be best used for first-passes over search output data before a more detailed manual inspection with graphical selection tools."
    },
    "1005/1005.1966_arXiv.txt": {
        "abstract": "We investigate the radio and $\\gamma$-ray beaming properties of normal and millisecond pulsars by selecting two samples from the known populations. The first, Sample G, contains pulsars which are detectable in blind searches of $\\gamma$-ray data from the \\emph{Fermi} Large Area Telescope. The second, Sample R, contains pulsars detectable in blind radio searches which have spin-down luminosities $\\dot{E} > 10^{34}$~erg s$^{-1}$. We analyse the fraction of the $\\gamma$-ray-selected Sample G which have detectable radio pulses and the fraction of the radio-selected Sample R which have detectable $\\gamma$-ray pulses. Twenty of our 35 Sample G pulsars have already observed radio pulses. This rules out low-altitude polar-cap beaming models if, as is currently believed, $\\gamma$-ray beams are generated in the outer magnetosphere and are very wide. We further find that, for the highest-$\\dot{E}$ pulsars, the radio and $\\gamma$-ray beams have comparable beaming factors, i.e., the beams cover similar regions of the sky as the star rotates. For lower-$\\dot{E}$ $\\gamma$-ray emitting pulsars, the radio beams have about half of the $\\gamma$-ray sky coverage. These results suggest that, for high-$\\dot{E}$ young and millisecond pulsars, the radio emission originates in wide beams from regions high in the pulsar magnetosphere, probably close to the null-charge surface and to the $\\gamma$-ray emitting regions. Furthermore, it suggests that for these high-$\\dot E$ pulsars, as in the $\\gamma$-ray case, features in the radio profile represent caustics in the emission beam pattern. ",
        "introduction": "The observed pulses from pulsars are interpreted as an emission beam, fixed relative to the rotating neutron star, sweeping across the Earth. The magnetospheric emission-region structure, the viewing angle with respect to the rotation axis and the angle between the magnetic and rotation axes define the shape of the observed pulse profile. Pulsar radio emission is generally thought to originate within a few percent of the light-cylinder radius ($R_{\\rm LC}$) above the polar caps \\citep[e.g.][]{lm88}.  Many pulsars exhibit a ``radius-to-frequency'' mapping \\citep{cor78} which is indicative of polar cap emission. The increasing pulse width at lower radio frequencies is interpreted as emission at increasing heights above the poles along more widely directed field lines.  The radio profiles of many young pulsars and millisecond pulsars (MSPs) have an ``interpulse'', a second pulse component separated from the main pulse by close to $180\\degr$ of pulse phase, or widely separated pulse components. In most cases, such pulsars do not exhibit radius-to-frequency mapping, that is, they have frequency-independent pulse component separations \\citep{hf86}.  The known population of $\\gamma$-ray pulsars has recently been greatly expanded by the six-month data release of the \\emph{Fermi} Large Area Telescope (LAT) \\citep[][hereafter A10]{aaa+10c}. This sample of 46 $\\gamma$-ray pulsars includes both young pulsars and MSPs, all with high spin-down luminosities: \\begin{equation} \\dot{E}=\\frac{4\\pi^{2}I\\dot{P}}{P^{3}}>10^{33}\\,{\\rm erg\\,s^{-1}}, \\end{equation} where $I=10^{45}$\\,gm\\,cm$^{-2}$ is the assumed neutron star moment of inertia, $P$ is the rotation period in seconds and $\\dot{P}$ is the period time-derivative.  Recent studies indicate that $\\gamma$-ray emission from pulsars originates at high altitudes in the pulsar magnetosphere, near $R_{\\rm LC}$ and from the last open field lines \\citep[e.g.,][A10]{aaa+09,aaa+09d}. Two classes of high-altitude particle acceleration and emission model are generally considered: slot-gap accelerators \\citep{mh04a,dr03} from the polar cap outwards along the last open field lines, and outer-gap accelerators \\citep{ry95,cz98} where the last closed field lines cross the null-charge surface defined by ${\\bf \\Omega\\cdot B} = 0$ ($\\bf \\Omega$ is the rotation velocity vector, and $\\bf B$ is the magnetic field vector). These models result in what we shall term \\emph{wide-beam} emission geometries.  Low-altitude models for $\\gamma$-ray emission have been largely ruled out by the generally wide and multi-component $\\gamma$-ray pulse profiles and the existence of pulsed $\\gamma$-ray emission at energies greater than a few GeV; at low altitudes such emission is precluded by electron-positron pair absorption. Furthermore, \\citet{rw09} find that outer-gap models are statistically preferred over slot-gap models based on fits to LAT pulse profiles for a sample of bright $\\gamma$-ray pulsars.  Four of the five brightest $\\gamma$-ray pulsars detected by the Energetic Gamma-Ray Experiment Telescope (EGRET) are also radio pulsars \\citep{tho08}. This, together with the similar radio and $\\gamma$-ray pulse morphology for many young pulsars and MSPs, led \\citet{man05} to suggest that, for these high-$\\dot{E}$ pulsars, the radio emission is also emitted in wide beams from the outer magnetosphere near $R_{\\rm LC}$. Within the framework of polar cap emission, \\citet{jw06} and \\citet[][hereafter KJ07]{kj07} proposed that the radio emission from young pulsars is emitted at intermediate altitudes of $\\sim1000$\\,km or $\\sim 0.2 R_{\\rm LC}$ in order to account for wide profiles of young pulsars.  In this letter, we identify samples of $\\gamma$-ray-selected and radio-selected pulsars and consider the pulsars in each sample that also emit radio and $\\gamma$-ray pulses respectively. We then discuss the implications of our analysis for the radio beams of young pulsars and MSPs. ",
        "conclusions": "By comparing the relative occurrence of radio and $\\gamma$-ray pulsed emission in $\\gamma$-ray-selected and radio-selected samples, we find that the radio beaming fraction $f_r$ appears to be close to unity for the highest-$\\dot{E}$ pulsars, decreasing to $\\sim0.5$ for lower-$\\dot{E}$ $\\gamma$-ray pulsars. Our estimated lower bound $f_r$, averaged over all $\\gamma$-ray-selected pulsars, of $\\sim 0.57$ is inconsistent with low-altitude polar-cap models for the radio emission but is consistent with KJ07's intermediate-altitude polar-cap model. However, the KJ07 model is marginally inconsistent with our result that $f_r \\sim 1$ for the highest-$\\dot{E}$ pulsars. With other evidence from radio and $\\gamma$-ray pulse morphologies, these results suggest that, for high-$\\dot{E}$ pulsars, the radio emission originates in wide beams from the vicinity of the null-charge surface, possibly with a slot-gap configuration, with profile components representing caustics in the emission pattern.  Improved statistics from further $\\gamma$-ray and radio discoveries will help to test our conclusions. Modelling of radio pulse profiles with slot-gap or other wide-beam emission geometries for both young high-$\\dot{E}$ pulsars and MSPs would be valuable. If outer-magnetoshpere wide-beam radio emission from these pulsars is confirmed, it would have profound implications for pulsar population and evolution studies as well as our understanding of radio and $\\gamma$-ray pulse emission mechanisms."
    },
    "1005/1005.3478_arXiv.txt": {
        "abstract": "We utilize detailed time-varying  models of the coupled evolution of stars  and the  HI, $\\rm  H_2$, and  CO-bright H$_2$  gas  phases in galaxy-sized  numerical  simulations  to  explore the  evolution  of gas-rich and/or  metal-poor systems, expected to be  numerous in the Early Universe.  The inclusion of the CO-bright H$_2$ gas phase, and the  realistic rendering  of  star formation  as an  H$_2$-regulated process (and  the new feedback  processes that this  entails) allows the most realistic tracking  of strongly evolving galaxies, and much better comparison  with observations.   We find that  while galaxies eventually settle into  states conforming to Schmidt-Kennicutt (S-K) relations,  significant  and  systematic  deviations of  their  star formation rates (SFRs) from  the latter occur, especially pronounced and prolonged  for metal-poor systems.  The  largest such deviations occur  for gas-rich  galaxies during  early evolutionary  stages but also  during brief  periods at  later stages.   Given  that gas-rich and/or metal-poor  states of present-epoch galaxies  are expected in the Early Universe while a  much larger number of mergers frequently resets non-isolated  systems to gas-rich states,  even brief periods of sustained deviations  of their SFRs from those  expected from S-K relations  may come  to  characterize significant  periods of  their stellar   mass  built-up.    This   indicates  potentially   serious limitations of (S-K)-type relations as reliable sub-grid elements of star formation physics in  simulations of structure formation in the Early Universe.  We anticipate  that galaxies with marked deviations from the S-K  relations will be found at  high redshifts as unbiased inventories  of total  gas mass  become possible  with ALMA  and the EVLA. ",
        "introduction": "}  Since it was first proposed  as a phenomenological relation linking HI gas  mass and  star  formation surface  density  in galaxies  (Schmidt 1959),  and subsequently  better constrained  and re-formulated  as to include  also   CO-bright  H$_2$  gas  (Kennicutt   1998,  2008),  the Schmidt-Kennicutt  (hereafter S-K) relation:  $\\rm \\Sigma_{SFR}\\propto [\\Sigma   (HI)]^k$   (k$\\sim   $1-2),   has  provided   the   standard observational framework relating the star formation rate (as a surface density rate $\\rm \\Sigma_{SFR}$) to the gas supply in galaxies.  It is also  an important  element  of the  sub-grid  star formation  physics incorporated in galaxy evolution  and structure formation models (e.g. Baugh et  al.  2005  and references therein;  Springel, Di  Matteo, \\& Hernquist  2005,  Schaye  \\&   Della  Vecchia  2008)  where  numerical resolution  limitations preclude  a  more detailed  treatment of  star formation over the scales involved.  Many theoretical (Dopita \\& Ryder 1994; Robertson  \\& Kravtsov 2008),  and observational (e.g.   Wong \\& Blitz 2002; Bigiel et al.  2008) studies have been made to demonstrate its  validity,  with the  most  important  recent  advances being  the identification of the CO-bright H$_2$ gas as better correlated to star formation in galaxies  than atomic hydrogen (Wong \\&  Blitz 2002), and the   direct   star-formation   role   of  its   dense   phase   ($\\rm n(H_2)>10^4\\,cm^{-3}$) with a k$\\sim $1 (Gao \\& Solomon~2004).   Unfortunately past analytical and  numerical investigations of the S-K relation did not include a  multi-phase ISM (though see Gerritsen 1997 for  an early  investigation  that includes  it)  or assumed  sub-grid models reacting instantaneously to changes  in the global state of the ISM. Thus they are ill-suited  to explore very gas-rich systems and/or early  galaxy evolutionary  stages, when  the various  ISM  phases and their  interplay  with  the   stellar  content  have  not  established equilibrium.   Finally, such  models  cannot be  compared directly  to observations since they do not include the H$_2$ gas phase (the direct fuel of  star formation),  or do so  but assume  that all such  gas is CO-bright (e.g.   Gnedin et  al. 2009).  The  latter is not  the case, especially in a metal-poor and/or far-UV intense ISM environment (e.g. Israel 1997; Maloney \\& Black 1988; Pak et al.  1998), which is common during  the  gas-rich  and  vigorously star-forming  epochs  in  early galaxy~evolution.  \\subsection{ISM+stars galaxy models: features and limitations }  Here, we use  our time-varying models of the  coupled evolution of HI, H$_2$ gas  phases and stars  in galaxy-sized numerical  simulations to investigate  the  emergence  and  possible  deviations  from  the  S-K relations.   Full details  and tests  of our  method can  be  found in Pelupessy,  Papadopoulos, \\&  van  der Werf  (2006)  and Pelupessy  \\& Papadopoulos (2009).  We use an N-body/SPH code and solve for the full thermodynamic evolution of  the WNM and CNM HI  phases (see Wolfire et al. 2003)  assuming neither an equillibrium nor  an Effective Equation of  State  (EOS), unlike  most  current  cosmological or  galaxy-sized structure formation  models (e.g.  Springel \\& Hernquist  2003; Cox et al.  2006a,b; Narayanan et al.   2009).  It is worth pointing out that in  such models  the  coldest ISM  phase  tracked is  usually at  $\\rm T_{kin}$$\\sim $10$^4$\\,K (i.e.  thermodynamically far removed from the one  truly  forming the  stars),  and  a Schmidt-Kennicutt  (S-K)-type relation  between that  phase and  star formation  rate  is postulated (e.g.   Kravtsov et  al.  2004;  Governato et  al.  2007).   In recent models  colder  gas  ($\\sim  $300\\,K) is  tracked,  but  instantaneous equillibrium gas  thermodynamics remains an assumption  (e.g Tasker \\& Bryan 2008) while the molecular  gas phase and the feedback effects of star formation on it are not  included (e.g.  Tasker \\& Tan 2009).  In our approach star formation is controlled by gravitational instability via a  Jeans mass criterion, the  thermal state of the  gas is tracked explicitely, while an H$_2$-richness  criterion for star formation can be  applied in  addition  to that  of  gravitational instability  (see Pelupessy \\& Papadopoulos 2009 for details).   The code tracks the H$_2$ phase with a physical model for substructure using a minimal set of  assumptions. It follows the H$_2$ formation on dust grains and its  thermal \\& far-UV induced destruction, accounting for self-shielding  and dust shielding.  The CO-bright  H$_2$ phase is indentified  as  a  post-processing  step, using  the  most  important chemical reactions  (R{\\\"o}llig at al.  2006) solving  for the C$^{+}$ envelope per gas  cloud, as regulated by the  local far-UV field.  The aspects  making  our  code  particularly suitable  for  examining  the validity of  the S-K  relation over  a range of  conditions are:  a) a time-varying  treatment  of   the  H$_2$$\\leftrightarrow$HI  gas  mass exchange  while tracking  the ISM  from Warm  Neutral Medium  (WNM) HI ($\\rm T_k$$\\sim  $10$^4$\\,K, $\\rm n\\sim  $(0.1--1)\\,cm$^{-3}$) to Cold Neutral Medium (CNM) HI and H$_2$ gas ($\\rm T_{k}$$\\sim $(30--200)\\,K, $\\rm n$$\\sim$(10--200)\\,$\\rm cm^{-3}$), b) the versatility of using an H$_2$-regulating star  formation criterion,  {\\it in addition}  to the regular gravitational  instability criterion, and c)  CO formation and destruction  (for this  specific  version of  our  code).  The  latter allows  direct comparisons  with observations  since CO  line emission rather than H$_2$ is the real observable in~galaxies.   The most  realistic models  employ our H$_2$-regulated  star formation which uses the  local $\\rm H_2$ gas mass fraction  as a star formation regulator  in  the  dynamical  setting  of  an  evolving  galaxy.   We implement this  molecular regulated (MR) star  formation by converting only  the molecular  ($\\rm  f_m$)  mass fraction  of  an unstable  gas particle  to stars.   Unlike  our other  simulations  (where H$_2$  is tracked  but plays  no  role in  star  formation) that  need an  adhoc parameter  $\\epsilon_{SF}$  designating the  local  gas mass  fraction converted into stars, the MR  models contain a physical basis for this parameter while retaining the original  requirement of SF gas as Jeans unstable  (see Pelupessy  et  al  2006 for  more  details and  tests). Finally here we  limit ourselves to systems with  $\\rm M_{baryons} \\le 10^{10}\\,M_{\\odot}$  where we can  maintain the  ability to  track the full  range  of  gas  densities  and temperatures  necessary  for  the HI$\\rightarrow $H$_2$ phase transition to take place (typically in the densest and coldest regions of CNM HI~gas). ",
        "conclusions": "Our  results  are of  special  importance  for  the modeling  of  very gas-rich and/or metal-poor progenitors of present-epoch galaxies found in the  distant Universe,  or systems where  major gas  mass accretion events frequently ``reset'' their evolutionary states back to gas-rich ones.   In  such cases  the  non-equilibrium, non-linear,  mass/energy exchange between the various ISM  phases and the stellar component may come  to dominate significant  periods of  intense star  formation and stellar mass built-up during which the S-K relation is not applicable. In short  we find that it  may work well  for present-epoch metal-rich spirals with  modest remaining gas  mass fractions (i.e.   systems for which the (S-K) relations were  originally deduced), but not for their very  gas-rich/metal-poor  progenitors  in  the Early  Universe.   The importance of such deviations of  the actual star formation rates from those  expected from  S-K  relations does  not  lie so  much in  their magnitude (though  this must be explored further  for gas-rich systems larger  than the  ones  modelled here)  rather  than their  systematic nature  (e.g.   star-forming  systems  spending certain  periods  with always higher or  always lower SFRs than those  estimated from the S-K relations).   It is  the latter  that may  make  such phenomenological relations poor  choices for the sub-grid physics  of star-formation in rigorous structure formation models in a cosmological~setting.    Finally we note that when it comes to the large gas-rich galaxies that are currently accessible at high redshifts our results, drawn for less massive  systems, remain provisional.   Nevertheless for  more massive and very  gas-rich star-forming systems  the larger amplitudes  of ISM equilibrium-perturbing  agents (e.g.   SNs, far-UV  radiation fields), and the shorter timescales  characterizing their variations, will more likely  than not  exaggerate  the deviations  of  true star  formation versus  the one expected  from (S-K)-type  phenomenological relations. An  unbiased  observational effort  to  find  and study  (S-K)-deviant galaxies at high redshifts (soon to be possible with ALMA and the EVLA over a  wide range  of galaxy masses),  as well as  extending detailed numerical  modeling  of  gas  and  stars towards  larger  systems  (as computational capabilities improve), are important in establishing the validity range of S-K relations and thus their utility as an important sub-grid element in models of structure formation in the Universe."
    },
    "1005/1005.5095.txt": {
        "abstract": "The study of multiple extrasolar planetary systems has the opportunity to obtain constraints for the planetary masses and orbital inclinations via the detection of mutual perturbations. The analysis of precise radial velocity measurements might reveal these planet-planet interactions and yields a more accurate view of such planetary systems. Like in the generic data modelling problems, a fit to radial velocity data series has a set of unknown parameters of which parametric derivatives have to be known by both the regression methods and the estimations for the uncertainties. In this paper an algorithm is described that aids the computation of such derivatives in case of when planetary perturbations are not neglected. The application of the algorithm is demonstrated on the planetary systems of HD~73526, HD~128311 and HD~155358. In addition to the functions related to radial velocity analysis, the actual implementation of the algorithm contains functions that computes spatial coordinates, velocities and barycentric coordinates for each planet. These functions aid the joint analysis of multiple transiting planetary systems, transit timing and/or duration variations or systems where the proper motion of the host star is also measured involving high precision astrometry. The practical implementation related to the above mentioned problems features functions that make these kind of investigations rather simple and effective. ",
        "introduction": "\\label{sec:introduction}  As of this writing, $36$ multiple planetary systems are known around main sequence stars. Most of these systems bear two detected planets, $8$ of them have $3$ and $2$ of them have $4$ planets while the star 55~Cnc has 5 companions\\footnote{See e.g. \\texttt{http://exoplanet.eu} for an up-to-date list.}. With the exception of HR~8799 \\citep{marois2008}, all of the detections are based on or confirmed by the measurements of the radial velocity (RV) variations of the host stars\\footnote{The planetary system around HR~8799 with 3 confirmed planets has been detected by direct imaging.}. The planet HAT-P-13b \\citep{bakos2009} also transits its host star. The detection of this planet was based on transit photometry while the second companion in this system has been revealed by radial velocity measurements. In general, analysis of RV variations constrains the mass ($m$) of planets by a lower limit. Namely, only the quantity $m\\sin i$ is determined by RV data where $i$ is the orbital inclination (relative to the tangential plane of sky). With the exception of transiting planets and systems were the spatial motion of the host star is detected via astrometry, there is no direct evidence for the actual value of the orbital inclination (and therefore the mass of the planet). In case of transiting systems, inclinations are constrained by measuring the impact parameter from the light curves \\citep[see e.g.][]{pal2010}, while astrometry yields not only the inclination but the orientation of the orbital plane as well \\citep{bean2009,benedict2010,mcarthur2010}. The planetary system around GJ~876 is the only known one where mutual inclination is also detected with a 2-$\\sigma$ confidence \\citep{bean2009}. A great advantage of multiple planetary systems is the possibility of detecting mutual perturbations via the deflection of RV values from the purely Keplerian solution \\citep[see e.g.][]{laughlin2001}. Therefore, precise analysis of accurate RV series may yield to an acceptable constraint for both the inclination and the planetary masses. Additionally, planet-planet interactions depend on the mutual inclinations, thus more complex models (such as non-coplanar orbits) for the whole planetary systems can be investigated.  Like in the majority of data modelling problems, RV variations (of the host star) in single or multiple planetary systems are modelled with a function of a few external (unknown) parameters. These parameters include the orbital elements and the masses of the planets as well as the barycentric velocity of the host star. For most of the regression methods involved in data modelling \\citep[see e.g. the Levenberg-Marquard algorithm,][]{press1992}, and for the analytic estimation of the covariances, uncertainties and correlations of the model parameters \\citep[see e.g.][]{finn1992,pal2009}, the partial derivatives of the model functions (with respect to the model parameters) have to be known in advance. The simplest way of RV curve modelling does not take into account the mutual interactions between the planets and characterizes the observed RV variations as a sum of independent Keplerian models. \\cite{wright2009b} describes an algorithm detailing the efficient computation of the parametric derivatives of RV model functions where the mutual planetary perturbations are neglected. The main objective of this paper is to present an algorithm that calculates parametric derivatives when the planet-planet interactions are also taken into account.  The structure of the paper is as follows. In the next section, we describe the mathematical tools used to construct the algorithm itself (including the discussion of optimal orbital parameterization, and the numerical integration). In Section~\\ref{sec:implementation} the practical implementation is detailed while in Section~\\ref{sec:applications} we demonstrate the usage of the algorithm for three specific multiple planetary systems. The results are summarized in the last section. ",
        "conclusions": "In this paper we described an algorithm based on the Lie-integration method that efficiently computes the parametric derivatives of radial velocity model functions for multiple planetary systems when the planet-planet interactions are also taken into account. The analysis of these systems yields more accurate constrains for planetary masses since the orbital and mutual inclinations can also be derived if precise radial velocity data are available. Additionally, the presented analytic formulae and integration method aid to plan observation schedules in order to optimize the telescope time utilization in order to detect planetary perturbations.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1005/1005.0637_arXiv.txt": {
        "abstract": "Gas-phase complex organic molecules have been detected toward a range of high- and low-mass star-forming regions at abundances which cannot be explained by any known gas-phase chemistry. Recent laboratory experiments show that UV irradiation of CH$_3$OH-rich ices may be an important mechanism for producing complex molecules and releasing them into the gas-phase. To test this ice formation scenario we mapped the B1-b dust core and nearby protostar in CH$_3$OH gas using the IRAM 30m telescope to identify locations of efficient non-thermal ice desorption. We find three CH$_3$OH abundance peaks tracing two outflows and a quiescent region on the side of the core facing the protostar. The CH$_3$OH gas has a rotational temperature of $\\sim$10~K at all locations. The quiescent CH$_3$OH abundance peak and one outflow position were searched for complex molecules. Narrow, 0.6-0.8 km s$^{-1}$ wide, HCOOCH$_3$ and CH$_3$CHO lines originating in cold gas are clearly detected, CH$_3$OCH$_3$ is tentatively detected and C$_2$H$_5$OH and HOCH$_2$CHO are undetected toward the quiescent core, while no complex molecular lines were found toward the outflow. The core abundances with respect to CH$_3$OH are $\\sim$2.3\\% and 1.1\\% for HCOOCH$_3$ and CH$_3$CHO, respectively, and the upper limits are 0.7--1.1\\%, which is similar to most other low-mass sources. The observed complex molecule characteristics toward B1-b and the pre-dominance of HCO-bearing species  suggest a cold ice (below 25~K, the sublimation temperature of CO) formation pathway followed by non-thermal desorption through e.g. UV photons traveling through outflow cavities. The observed complex gas composition together with the lack of any evidence of warm gas-phase chemistry provide clear evidence of efficient complex molecule formation in cold interstellar ices. ",
        "introduction": "Complex organic molecules have been detected toward a range of astrophysical environments, including low-mass protostars \\citep{vanDishoeck95, Cazaux03, Bottinelli07}; however, the origins of these complex molecules as well as their fates are uncertain. Commonly-suggested formation routes for the detected molecules include various gas-phase reactions starting with thermally evaporated CH$_3$OH ice, atom-addition reactions on dust grains, and UV- and cosmic ray-induced chemistry in the icy grain mantles that form during the pre-stellar stages \\citep[][for a review]{Charnley92, Nomura04, Herbst09}. The focus is currently on an ice formation pathway \\citep[e.g.][]{Garrod06, Garrod08} because of the failures of gas phase chemistry to explain the observed abundances of some of the most common complex molecules toward low-mass protostars, especially HCOOCH$_3$.  Recent experiments on the photochemistry of CH$_3$OH-rich ices have shown that 1) UV irradiation of CH$_3$OH ices at 20--70~K result in the production of large amounts of the complex molecules observed around protostars and 2) the chemistry has a product branching ratio which is temperature and ice composition dependent \\citep{Oberg09d}. A key result is that HCOOCH$_3$ and other HCO-X ices are only abundantly produced in CO:CH$_3$OH mixtures. A similar result has been reported for proton-bombarded CH$_3$OH and CH$_3$OH:CO ices \\citep{Bennett07b}. CO:CH$_3$OH ice mixtures are probably common since CH$_3$OH forms from hydrogenation of CO \\citep[e.g.][]{Watanabe03,Cuppen09}. CO evaporates at 17--25~K on astrophysical time scales \\citep{Bisschop06}, and thus HCO-X ices will mainly form in cold, UV-exposed ices. At higher temperatures, closer to the protostar, UV irradiation of the remaining pure CH$_3$OH ice will instead favor the production of C$_2$H$_5$OH, CH$_3$OCH$_3$ and (CH$_2$OH)$_2$.  In light of this proposed formation scenario, protostars rich in CH$_3$OH ice are natural targets when searching for complex molecule sources. The low-mass protostar B1-b is such a source. From previous SCUBA and 3 mm continuum maps, B1-b consists of two cores B1-bN and B1-bS, separated by 20'' \\citep{Hirano99}. The cores are similar with $T_{\\rm dust}\\sim18$~K, $M=1.6-1.8\\:M_{\\odot}$ and $L_{\\rm bol}=2.6-3.1\\:L_{\\odot}$. The {\\it Spitzer Space Telescope} only observed one protostar, to the south-west of the B1-bS dust core at 03:33:20.34, +31:07:21.4 (J2000), but it was still named B1-b -- to avoid confusion it will be referred to as the `protostar' position. From the {\\it Spitzer} $c2d$ (From Molecular Cores to Planet Forming Disks) ice survey the B1-b protostar has a CH$_3$OH ice abundance of 11\\% with respect to H$_2$O ice \\citep{Boogert08}, corresponding to $\\sim5\\times10^{-6}$ with respect to H$_2$. The fractional abundance of H$^{13}$CO$^+$ decreases by a factor of 5 toward the SCUBA core, indicative of CO freeze-out \\citep{Hirano99}. A HCO-dominated complex chemistry is thus expected where the ice is exposed to UV radiation.  The B1-b region is complicated by a number of outflows, which may enhance the UV field through shocks and open cavities through which stellar UV  photons can escape \\citep{Jorgensen06, Walawender09, Hiramatsu10}. The outflow of most interest to this study runs in the south-west direction from B1-b protostar, though it is not known whether B1-b is its source or just happens to lie in its path. The outflow is not observed north-east of the B1-b protostar, where the B1-b SCUBA core is situated. It is therefore unclear whether the outflow terminates at the protostar or continues into the core, hidden from view. If it does penetrate into the SCUBA core, this may increase the UV flux in the region by orders of magnitude, enhancing both the UV photochemistry and photodesorption of ices.  The latter is important, since once formed, the complex molecules must be (partly) released into the gas phase to be observable at millimeter wavelengths. The second reason for targeting B1-b is the detection of a large CH$_3$OH gas phase column density of $\\sim$2.3--2.5$\\times10^{14}$ cm$^{-2}$ toward the protostar \\citep{Oberg09a,Hiramatsu10} and thus ice evaporation. The narrow line width and low excitation temperature of the observed CH$_3$OH suggest non-thermal ice evaporation -- the abundances are consistent with UV photodesorption. Assuming that complex ice chemistry products desorb non-thermally at an efficiency similar to CH$_3$OH ($\\sim$10$^{-3}$ molecules per incident UV photon \\citep{Oberg09d}),  the protostellar environment should contain a cold gas-phase fingerprint of the complex ice composition; B1-b may offer the first possibility to observe the earliest stages of complex molecule formation, untainted by either warm ice chemistry or gas-phase processing.  Addressing these predictions, we present IRAM 30-m telescope observations of a $80''\\times80''$ CH$_3$OH map of the B1-b protostellar envelope and nearby dust cores, followed by a search for HCOOCH$_3$, CH$_3$CHO, CH$_3$OCH$_3$, C$_2$H$_5$OH and HOCH$_2$CHO toward two of the identified CH$_3$OH peaks. The CH$_3$OH gas abundances and the complex abundances and upper limits with respect to CH$_3$OH are then discussed in terms of different ice desorption mechanisms and complex molecule formation scenarios. The results are finally compared with complex molecule observations toward other low-mass protostars and outflows. ",
        "conclusions": "\\begin{enumerate} \\item Cold CH$_3$OH gas (excitation temperature of $\\sim$10~K) is abundant and widespread toward B1-b. Quiescent CH$_3$OH is most abundant in between the B1-b SCUBA dust core and the B1-b protostar, indicative of UV photodesorption of ice on the side of the quiescent dust core because of radiation from the protostar escaping through outflow cavities. \\item Cold HCOOCH$_3$ and CH$_3$CHO are both detected and CH$_3$OCH$_3$ is tentatively detected toward the quiescent CH$_3$OH peak. No complex molecules are observed toward an outflow associated with B1-b, but 3-$\\sigma$ upper limits with respect to CH$_3$OH are only slightly lower compared to the quiescent core. In addition an asymmetry in the CH$_3$CHO emission lines is suggestive of some contribution from a small outflow included in the beam. \\item Assuming a 10~K excitation temperature the calculated beam-averaged abundances are 2.3\\% for HCOOCH$_3$, 1.1\\% for CH$_3$CHO, $\\lesssim$0.8\\% for CH$_3$OCH$_3$ and $<$1.0-1.1\\% for C$_2$H$_5$OH and HOCH$_2$CHO with respect to CH$_3$OH. \\item Building on recent experiments, the observations of large abundances of HCO-containing molecules are explained by UV/cosmic ray processing of cold CH$_3$OH:CO ice followed by non-thermal desorption of a fraction of the produced organic ice. \\item The beam-averaged abundances with respect to CH$_3$OH and the ratios between different complex molecules are similar to other single-dish observations toward low-mass protostars. In contrast resolved observations of the cores around protostars are more abundandant in complex molecules without a HCO-group. This is consistent with a complex ice chemistry that evolves from HCO-rich to HCO-poor as the ice warms up and the CO-ice evaporates when the icy dust falls toward  the protostar. \\end{enumerate}  Including B1-b, there are still only a handful of observations of complex molecules toward low-mass protostars. In addition few complex molecules are detected toward each source. This makes it difficult both to determine the prevalence of a complex organic chemistry around low-mass protostars and to firmly establish the main formation path of such molecules. The detections of complex molecules in the vicinity of the CH$_3$OH-ice-rich protostar suggests that targeting other CH$_3$OH-ice rich protostars may increase our sample of complex molecule detections significantly. Maybe more important from a formation pathway point of view is to increase our understanding on how the complex chemistry varies between warm, luke-warm and cold regions i.e. how it varies with distance from the protostar. This should be pursued both on larger scales with single dish observations and on smaller scales with interferometers -- similarly to what has been done toward IRAS 16293 and NGC1333 2a \\citep[e.g.][]{Jorgensen05,Bisschop08}. In the meantime the complex molecules observed toward B1-b, and their exceptionally low excitation temperatures and line widths, provides clear evidence, for the efficient formation of complex ices during star- and planet-formation."
    },
    "1005/1005.5012_arXiv.txt": {
        "abstract": "{The physical conditions of the solar photosphere change on very small spatial scales both horizontally and vertically. Such a complexity may pose a serious obstacle to the accurate determination of solar magnetic fields.}  {We examine the applicability of Milne-Eddington (ME) inversions to high spatial resolution observations of the quiet Sun. Our aim is to understand the connection between the ME inferences and the actual stratifications of the atmospheric parameters.}  {We use magnetoconvection simulations of the solar surface to synthesize asymmetric Stokes profiles such as those observed in the quiet Sun. We then invert the profiles with the ME approximation. We perform an empirical analysis of the heights of formation of ME measurements and analyze the uncertainties brought about by the ME approximation. We also investigate the quality of the fits and their relationship with the model stratifications.} {The atmospheric parameters derived from ME inversions of high-spatial resolution profiles are reasonably accurate and can be used for statistical analyses of solar magnetic fields, even if the fit is not always good. We also show that the ME inferences cannot be assigned to a specific atmospheric layer: different parameters sample different ranges of optical depths, and even the same parameter may trace different layers depending on the physical conditions of the atmosphere. Despite this variability, ME inversions tend to probe deeper layers in granules as compared with intergranular lanes. }  {} ",
        "introduction": "\\label{cap5:sec:intro}  The solar spectrum carries information about the properties of our star. In general, a broad range of atmospheric layers contribute to the shape of the spectral lines, making it difficult to extract this information directly. Both the measurement process and the method of analysis introduce uncertainties in the physical quantities retrieved from the observations. Sources of error are photon noise and instrumental effects like limited spectral resolution, wavelength sampling, and angular resolution, but also the simplifications and approximations of the model used to interpret the measurements.   In this paper we want to evaluate the merits of Milne-Eddington (ME) inversions for the analysis of the polarization line profiles emerging from the solar atmosphere. The ME approximation does not account for vertical variations of the parameters \\citep{1956PASJ....8..108U, Rakk1,Rakk2}, so it cannot accurately describe the solar plasma when rapid changes in height are present.  What is, then, the significance of the ME parameters?  To answer this question it is necessary to simulate the processes of line formation and data inversion. Usually one prescribes a set of model atmospheres, performs spectral synthesis calculations, inverts the synthetic profiles, and compares the results with the known input. A common approach is to use ME models both to generate the spectra and to invert them \\citep[e.g.,][]{nortonetal2006, borreroetal2007}. In that case the analysis is internally consistent and the uncertainties of the retrieved ME parameters are mostly due to the noise and, to a smaller extent, to the convergence of the algorithm, provided that the spectral resolution and wavelength sampling are appropriate. Uncertainties caused by photon noise are known as statistical errors and can be evaluated by means of numerical tests or, more efficiently, by using ME response functions \\citep{2007A&A...462.1137O,2010ApJ...711..312D}. However, they represent only a small fraction of the total error. Another source of error is the very assumption of height-independent parameters, which leads to symmetric line profiles. What happens when realistic (i.e., asymmetric) Stokes spectra are analyzed in terms of ME models? Do the uncertainties of the retrieved parameters increase significantly? Answering these questions is the aim of the present work.  \\begin{figure*}[!t] \\centering \\resizebox{0.425\\hsize}{!}{\\includegraphics{Fig1a.ps}} \\resizebox{0.425\\hsize}{!}{\\includegraphics{Fig1b.ps}} \\resizebox{0.425\\hsize}{!}{\\includegraphics{Fig1c.ps}} \\resizebox{0.425\\hsize}{!}{\\includegraphics{Fig1d.ps}} \\caption{Magnetic field strength, inclination, azimuth, and LOS velocity at $\\tau=1$ in a simulation snapshot with $\\langle B \\rangle=140$~G. Negative velocities represent upflows.} \\label{Fig1} \\end{figure*}  A first study of the capabilities and limitations of ME inversions was carried out by \\cite{1998ApJ...494..453W} using simple (non-ME) model atmospheres. They made a quantitative comparison of results obtained with the ME code of the High Altitude Observatory \\citep{lites2,lites1} and the SIR code \\citep[Stokes Inversion based on Response functions;][]{1992ApJ...398..375R}. The main conclusion of their work was that ME inversions provide accurate values of the physical parameters averaged along the line of sight, at least when the stratifications are smooth.  More recently, \\cite{2007MmSAI..78..166K} have investigated whether the magnetic field stratification itself can be determined reliably through inversion of high resolution data. To that end, they synthesized the Stokes profiles of the \\ion{Fe}{i} 630~nm lines with the help of MHD models and inverted them with SIR, allowing for vertical gradients of the atmospheric parameters. The analysis showed that SIR is able to recover the actual magnetic stratification for fields as weak as 50~G if no noise is present. This work extends the results of \\cite{1998ApJ...494..453W} to the case in which the stratifications are not smooth.   To determine the uncertainties associated with ME inversions of asymmetric Stokes profiles we use state-of-the-art magnetohydrodynamic simulations (Sect.~\\ref{MHD}). Our goal is to describe the solar photosphere as realistically as possible. We construct model atmopheres from the simulations and synthesize the Stokes profiles of the \\ion{Fe}{i} 630.2~nm lines emerging from them (Sect.~\\ref{synthesis}). The SIR code is used for the spectral synthesis, so the profiles are asymmetric. Finally, we apply a ME inversion to the data (Sect.~\\ref{inversion}). In our numerical experiments, the spatial sampling of the MHD models, $0\\farcs 0287$, is preserved. There are two reasons why we neglect the effects of solar instrumentation: first, they have already been studied in the past (e.g., Orozco Su\\'arez et al. 2007, 2010); second, this sampling is close to critical for the observations to be delivered by large telescopes like the Advanced Technology Solar Telescope\\footnote{http://atst.nso.edu/} \\citep{wagner2006} and the European Solar Telescope\\footnote{http://www.iac.es/project/EST/} \\citep{collados2008}. We invert the profiles with the MILOS code \\citep{2007A&A...462.1137O}.\\footnote{MILOS is programmed in IDL and can be downloaded from our website, http://spg.iaa.es/download.asp} A direct comparison of the retrieved and true parameters allows us to determine the effective ``heights of formation'' of the ME parameters (Sect.~\\ref{results}) and to quantify the errors caused by the ME approximation (Sect.~\\ref{results2}). The conclusions of our work are given in Sect.~\\ref{conclus}.  For completeness, the results of ME inversions are compared with those of tachogram/magnetogram-like analyses in the Appendix. ",
        "conclusions": "\\label{conclus}  In this paper we have analyzed radiative MHD simulations of the quiet Sun. We have used them to synthesize Stokes line profiles in three different spectral regions (525.0, 617.3, and 630.2~nm). The comparison of the synthetic profiles with the FTS atlas suggests that the simulations describe quite satisfactorily the physical conditions of the solar photosphere, although the MHD models are slightly hotter than the HSRA around $\\tau=1$.  After synthesizing the Stokes profiles, the applicability of ME inversions to high spatial resolution observations has been examined. We have considered the case of the \\ion{Fe}{i} lines at 630.2~nm. The analysis of the profiles by means of ME inversions has allowed us to characterize the uncertainties that can be expected from the ME approximation. For this reason, the synthetic profiles were not degraded by noise, instrumental effects, or spatial resolution.  The main limitation of ME inversions is that they provide constant atmospheric parameters, whereas the MHD models feature physical properties that change with height. This limitation means that ME models are unable to reproduce spectral line asymmetries. Consequently, the ME inferences cannot be assigned to a specific optical depth. Depending on the conditions of the atmosphere, the retrieved ME parameters sample different layers.  However, from a statistical point of view we conclude that ME inversions provide fair estimates of the physical conditions prevailing at $\\log\\tau \\sim -1$. The rms uncertainty is smaller than 30~G for the magnetic field strength, 13\\degree\\/ and 20\\degree\\/ for the field inclination and azimuth, and 500~\\ms\\/ for the LOS velocity. Thus, ME inversions are appropriate for statistical analyses of the solar photosphere. This being said, it is important to realize that the errors may be large for individual pixels, even if the best-fit profiles reproduce the observations satisfactorily (the field strength in case {\\em a} of Fig.~\\ref{Fig5} is a good example of this).  Finally, we want to stress that the uncertainties associated with the ME approximation are larger than those due to photon noise \\citep{2006ASPC..358..197O,2010ApJ...711..312D}. However, the noise has another undesirable effect: it hides the weaker polarization signals. This fact has not been considered in our study.  \\begin{figure*}[!t] \\centering % \\resizebox{0.98\\hsize}{!}{\\includegraphics{FigApp.eps}} \\caption{Scatter plots of the longitudinal magnetic field, $B_{z}$ (first two rows), of the LOS velocity (third row), and of the magnetic field inclination (fourth row), as functions of their corresponding values in the MHD simulations at $\\log\\tau = -1$. The first column refers to results obtained with the 630.15 nm line; the second column displays results obtained with the 630.25 nm line; results with the two lines at the same time are shown in the third column. Color and line codes are the same as in Fig.\\ \\ref{Fig8}. The mean value of the rms difference between the inferences and the simulations at $\\log\\tau = -1$ is given in the upper left corner of each panel.} \\label{FigApp} \\end{figure*}  \\appendix"
    },
    "1005/1005.5362_arXiv.txt": {
        "abstract": "The bright and highly variable X-ray and radio source known as Cygnus X-3 was among the first X-ray sources discovered, yet it remains in many ways an enigma. Its known to consist of a massive, Wolf-Rayet primary in an extremely tight orbit with a compact object. Yet one of the most basic  of parameters - the mass of the compact object - is not known. Nor is it  even clear whether its is a neutron star or a black hole. In this Paper we present our analysis of the broad-band high-energy continua covering a substantial  range in luminosity and spectral morphology. We apply these results to a recently identified scaling relationship which has been demonstrated to provide reliable estimates of the  compact object mass in a number of accretion powered binaries.  This analysis leads us to conclude that the compact object in Cygnus X-3 has a mass greater than $4.2M_\\odot$ thus clearly indicative of a black hole and as such resolving a long-standing issue. The full range of uncertainty in our analysis and from using a range of recently published distance estimates constrains the compact object mass to lie between $4.2M_\\odot$ and $14.4M_\\odot$. Our favored estimate, based on a 9.0 kpc distance estimate is $\\sim 10 M_\\odot$ with the error margin of 3.2 solar masses. This result may thus pose challenges to shared-envelope evolutionary models of compact binaries, as well as establishing Cygnus X-3 as the first confirmed accretion-powered galactic gamma-ray source. ",
        "introduction": "The bright X-ray binary source known as Cygnus X-3 (herein Cyg X-3) was discovered over 40 years ago  by \\cite{gia67}, and has been the subsequent focus of extensive study. It nonetheless remains enigmatic in that the nature of the compact object; a neutron star or black hole has not been unambiguously determined, its high-amplitude intensity and spectral variations in the X-ray and radio bands are not well understood, and its low-amplitude and featureless power-density spectrum does not resemble other known neutron star or black hole X-ray binaries, e.g. \\cite{McC99}. Many of these unique attributes are due to the nature of the donor star, widely believed to be a Wolf Rayet (WR) star, e.g. \\cite{vankerk92} . It is also the strongest radio source among X-ray binaries, with a quiescent flux $<100$ mJy, a flux of 0.1 Jy flaring to several Jy in high-activity states. It has recently been detected in the ~0.1-1-GeV domain with the Fermi Gamma-Ray Space Telescope \\citep{abdo09} and is thus among a small number of gamma-ray bright X-ray binaries. Its  orbital period of 4.8 hours is typical for a low-mass binary, but its WR donor star could have a mass upwards of 30 M$_\\odot$, thus making it an extremely tight binary system, all the more so if the compact star is a black hole.  Distance estimates to Cyg X-3 range from 7.2 to  9.3 kpc; see details in \\cite{lzt09}. Here we adopt the distance of 9 kpc, inferred from dust-halo scattering measurements \\citep{pred00} however, we consider the full range of these distance determinations in evaluating the uncertainty of our result. Such a large distance means that it is viewed thorough many magnitudes of visual extinction. It is in addition likely to be enshrouded in the dense wind environment of the WR donor star. This combination of high local and Galactic line-of-sight column densities has rendered the usual photometric and spectroscopic techniques for deriving a binary solution intractable, and as a result, no reliable mass estimate for the compact has emerged. The local absorption has also limited conclusions based on X-ray spectroscopy, which may be further complicated by the presence of a synchrotron component associated with the radio and gamma-ray emission and the effects of Compton downscattering. While hard-to-soft spectral transitions suggestive of a black hole binary are evident, e.g. \\cite{hz09} they cannot always be easily reconciled within the context of familiar black-hole low-hard to high-soft state transition patterns, nor do they resemble neutron-star Z or atoll source spectral behavior. No pulsations are detected, which would be a clear signature of a magnetized neutron star, however, it is possible that pulsed (radio or X-ray) emission could be present in the system but unobservable due to the effects of dense ambient plasma. It is noted that this type of scenario has been proposed and is considered credible in explanation of another radio-loud, gamma-ray-bright X-ray binary LSI +61 303.  Previous  estimates of the compact object mass based on a variety of methods have been attempted,  but the results tend to be highly uncertain.  For example Schmutz, Geballe and Schild (1996) employing IR spectroscopy obtained a likely mass in the range of $7-40 M_\\odot$ with a best estimate of $17 M_\\odot$. Hanson, Still and Fender (2000) constructed a radial velocity curve from which derive a mass function of $0.027 M_\\odot$, which leads to a blackhole mass of less than $10 M_\\odot$ but could still accommodate a neutron star depending on the true value of the binary inclination. Vilhu et al. (2009) assess orbital modulation of X-ray emission lines and find a most likely compact object masses between 2 and $8 M_\\odot$ while Hjalmarsditter et al. (2008) estimate a $\\sim 30 M_\\odot$ black hole based on interpretation of color-luminosity diagrams.  Our approach involves a multi-epoch spectral analysis and assessment of the inferred parameters. \\cite{st09} and \\cite{ts09}, hereafter  ST09 and TS09 respectively established well defined correlations between certain temporal and spectral parameters. Furthermore, these correlation curves observed in different sources exhibited differing patterns for different binaries, but that the variety of those patterns was limited, and thus scalable. Recently, the application of this method was extended to study of an another class of X-ray source, the ULX NGC 5408 X-1 (Strohmayer and Mushotzky 2009). That study resulted in tangible evidence for the existence of an intermediate mass object in that system.  The ubiquitous nature of  these correlations led to the suggestion that the underlying physical processes leading to the observed variability properties are closely tied to  the Comptonizing media. Furthermore, they vary in a well defined  manner as the source makes a transition between spectral states.  The fact that the same correlations are seen in many sources, which vary widely in both luminosity (and thus presumably with mass accretion rate) and spectral state, suggests a common set of underlying physical conditions. TS09 showed that in GRS 1915+105 the photon index monotonically increases with  disk mass accretion rate, followed by saturation. The radio luminosity does not correlate with low frequency quasi-periodic oscillations or X-ray luminosity for the entire range of spectral states, from low-hard to high-soft through intermediate.  In this Paper we present the results of a broad-band multi-epoch spectral analysis of Cyg X-3, applied to a novel approach for compact object mass determination, using photon index$-$disk mass accretion rate correlation first implemented by  ST09.   We compiled a data base of some 35 observations of Cyg X-3, primarily from the RXTE satellite, applied similar modeling methods, and studied the inferred parameter interdependencies. From the results of that analysis, we have applied the scaling laws derived in ST09 deriving an estimate of the compact object mass, which we find to be consistent with a black hole.  In \\S 2 we briefly discuss our methodology for mass-determination, which is described in detail elsewhere. In \\S 3 we describe our database compilation and analysis techniques, in \\S 4 we present  our results, draw comparison to similar results previously calibrated from objects with independent mass estimates, and in \\S 5 we summarize our conclusions. ",
        "conclusions": "We have studied a large set of data covering a number of spectral-state transitions and intensity variations in the galactic binary Cyg X-3. We find that for about half of the 35 ~2-100-keV spectra analyzed the observations are well represented by our modeling approach - an absorbed BMC model including a iron line and edge components and a high-energy cutoff term. We then examined the correlation between the inferred photon indices of the Comptonized spectral component and the mass accretion rate, a proxy for which is the normalization term of our model. Based on those results we have applied the scaling method employing previously determined results for the well known Galactic binary GRS 1915+105 as a reference dataset. The calibration of our method for GRS 1915+015 is supported by the dynamical mass determination of \\cite{gre01}. This analysis led us to a lower-limit mass determination of about $ 4.2M_\\odot$ and a best-estimate value of $\\sim 10.1\\pm3.2 M_\\odot$ assuming a 9.0 kpc distance. The previous success of this scaling method for the BH  mass determination strongly supports our results for Cyg X-3.  Another result of our study is that a basic prediction of the theory of the converging inflow is supported by the observations of the index-mass accretion rate correlation seen in various black-hole X-ray binaries. Specifically, we argue that the spectral-index vs $\\dot m$ flattening, or saturation, seen is an observational signature of the presence of a converging inflow. The implication is that {\\it this effect provides robust observational evidence for the presence of a black hole in Cyg X-3.}"
    },
    "1005/1005.2297_arXiv.txt": {
        "abstract": "{} {Within the framework of the \\texttt{HERM33ES} Key Project, using the high resolution and sensitivity of the \\textit{Herschel} photometric data, we study the compact emission in the Local Group spiral galaxy M\\,33 to investigate the nature of the compact SPIRE emission sources. We extracted a catalogue of sources at 250\\,\\micron\\ in order to investigate the nature of this compact emission. Taking advantage of the unprecedented \\textit{Herschel} resolution at these wavelengths, we also focus on a more precise study of some striking H$\\alpha$ shells in the northern part of the galaxy.} {We present a catalogue of 159 compact emission sources in M\\,33 identified by SExtractor in the 250\\,\\micron\\ SPIRE band that is the one that provides the best spatial resolution. We also measured fluxes at 24\\,\\micron\\ and H$\\alpha$ for those 159 extracted sources. The morphological study of the shells also benefits from a multiwavelength approach including H$\\alpha$, far-ultraviolet from \\textit{GALEX}, and infrared from both \\textit{Spitzer} IRAC 8\\,\\micron\\ and MIPS 24\\,\\micron\\ in order to make comparisons.} {For the 159 compact sources selected at 250\\,\\micron, we find a very strong Pearson correlation coefficient with the MIPS 24\\,\\micron\\ emission ($r_{24} = 0.94$) and a rather strong correlation with the H$\\alpha$ emission, although with more scatter ($r_{\\rm H\\alpha} = 0.83$). The morphological study of the H$\\alpha$ shells shows a displacement between far-ultraviolet, H$\\alpha$, and the SPIRE bands. The cool dust emission from SPIRE clearly delineates the H$\\alpha$ shell structures.} {The very strong link between the 250\\,\\micron\\ compact emission and the 24\\,\\micron\\ and H$\\alpha$ emissions, by recovering the star formation rate from standard recipes for \\HII\\ regions, allows us to provide star formation rate calibrations based on the 250\\,\\micron\\ compact emission alone. The different locations of the H$\\alpha$ and far-ultraviolet emissions with respect to the SPIRE cool dust emission leads to a dynamical age of a few Myr for the H$\\alpha$ shells and the associated cool dust.} ",
        "introduction": "% Within the framework of the open time key project ``{\\it Herschel} M\\,33 extended survey ({\\tt HERM33ES})'', we are studying the galaxy M\\,33 to understand the origin of various diagnostic lines in heating and cooling and other processes in the interstellar medium (ISM). We refer the reader to \\citet{2010A&A...501L...1K} for more details about the overall {\\tt HERM33ES} project goals, as well as for a first presentation of the PACS and SPIRE maps of the entire galaxy, together with spatially averaged spectral energy distributions.  The local group, late-type, spiral galaxy --M\\,33-- is an ideal target for studying the detailed processes of star formation (SF). Indeed, the distance, 840\\,kpc \\citep{1991ApJ...372..455F}, and inclination, 56 degrees \\citep{1994ApJ...434..536R}, of M\\,33 allow us to reach an unprecedented spatial resolution in the far-infrared of an external spiral galaxy and investigate the phenomena in the interstellar medium which lead to the SF. With {\\it Herschel} \\citep{2010A&A...001L...1P}, we are able to resolve 85\\,pc, 109\\,pc, and 170\\,pc structures of the cool dust in the SPIRE 250, 350, and 500\\,\\micron\\ bands, respectively, which is the typical spatial scale of giant \\HII\\ regions and their complexes. In this article, we follow a two-fold approach in order to study the nature of the compact 250\\,\\micron\\ emission and take advantage of the resolution and sensitivity reached by {\\it Herschel} to resolve the cool dust emission associated with \\HII\\ regions. ",
        "conclusions": "For the Local Group spiral galaxy M\\,33, using the unprecedented resolution and sensitivity of the {\\it Herschel} SPIRE photometric data, we focus on the compact emission and conclude the following.  \\begin{itemize} \\item We created a catalogue of 159 sources from the SPIRE 250\\,\\micron\\ emission using a high detection threshold to focus on the compact emission of the cool dust; the compact emission sources associated to the cool dust follow the spiral pattern of the galaxy. \\item Very high Pearson correlation coefficients, $r_{24} = 0.94$ for the 24\\,\\micron\\ emission and $r_{\\rm H\\alpha} = 0.83$ for the H$\\alpha$ emission, confirm that the 159 sources are closely linked to SF. \\item By using standard (H$\\alpha$+24\\,\\micron\\ and 24\\,\\micron\\ alone) SFR recipes for \\HII\\ regions provided by \\citet{2007ApJ...666..870C,2010ApJ...714.1256C}, we are able to calibrate the 250\\,\\micron\\ compact emission as an SFR tracer and provide conversion factors that can be used in further studies (Eqs.~\\ref{eq:sfrMix} and \\ref{eq:sfr24}). \\item The morphological study of a set of three H$\\alpha$ shells shows that there is a displacement between far-ultraviolet and the SPIRE bands, while the H$\\alpha$ structure is in general much more coincident with the cool dust. \\item The different locations of the H$\\alpha$ and far-ultraviolet emissions with respect to the SPIRE cool dust emissions leads to a dynamical age of a few Myr for a set of H$\\alpha$ shells and the associated cool dust. \\end{itemize}  \\vspace{-.4cm}"
    },
    "1005/1005.4687_arXiv.txt": {
        "abstract": "We present an analysis of the host galaxy dependencies of Type Ia Supernovae (SNe Ia) from the full three year sample of the SDSS-II Supernova Survey. We re-discover, to high significance, the strong correlation between host galaxy type and the width of the observed SN light curve, i.e., fainter, quickly declining SNe Ia favor passive host galaxies, while brighter, slowly declining Ia's favor star-forming galaxies. We also find evidence (at between 2 to $3\\sigma$) that SNe Ia are $\\simeq0.1\\pm0.04$ magnitudes brighter in passive host galaxies, than in star--forming hosts, after the SN Ia light curves have been standardized using the light curve shape and color variations: This difference in brightness is present in both the SALT2 and MCLS2k2 light curve fitting methodologies. We see evidence for differences in the SN Ia color relationship between passive and star--forming host galaxies, e.g., for the MLCS2k2 technique, we see that SNe Ia in passive hosts favor a dust law of $R_V = 1.0\\pm0.2$, while SNe Ia in star-forming hosts require $R_V = 1.8^{+0.2}_{-0.4}$. The significance of these trends depends on the range of SN colors considered. We demonstrate that these effects can be parameterized using the stellar mass of the host galaxy (with a confidence of $>4\\sigma$) and including this extra parameter provides a better statistical fit to our data. Our results suggest that future cosmological analyses of SN Ia samples should include host galaxy information. ",
        "introduction": "Over the last decade, Type Ia Supernovae (SNe Ia) have become important cosmological probes as they can be used to measure distances to high redshift ($z\\lesssim1.5$). In recent years, numerous samples of SN Ia have been compiled, e.g., CSP \\citep{2006PASP..118....2H}, SNLS \\citep{2006A&A...447...31A}, ESSENCE \\citep{2007ApJ...666..694W}, SDSS \\citep{2008AJ....135..338F}, CfA \\citep{2009ApJ...700.1097H}, and combined, we are approaching $\\sim1000$ spectroscopically-confirmed SNe Ia available for cosmological analysis \\citep{2010arXiv1004.1711A}. With such large samples, it is becoming increasingly important to understand the systematic uncertainties (photometric calibration, SN color variations, etc.) associated with using SNe Ia for cosmology, including any additional physical parameters that could reduce the intrinsic scatter of the population.  One such parameter could be related to the environment of the supernova. First and foremost, one would expect differences in the colors of SNe Ia based on the different dust content of their hosts, i.e., potential variations in local circumstellar dust around the progenitor star \\citep{wang2005, 2008ApJ...686L.103G} and/or differences in the global dust content of different galaxy types\\footnote{Dust in our own Galaxy is usually corrected for using the dust maps of \\citet{1998ApJ...500..525S}}.  Despite these concerns, most analyses account for dust during the fitting of the supernova light curves by assuming a single absorption law ($R_V$) for all SNe, which minimizes the scatter around the Hubble diagram \\citep{1998A&A...331..815T}. This process, however, has led to a dust law that is significantly different from the canonical value for our Galaxy ($R_V \\approx 3.1$), e.g., \\citet{2007ApJ...664L..13C} find $R_V \\approx 1$ for nearby SNe Ia, while  \\citet{2009ApJS..185...32K} obtained a best fit of $R_V = 2.18 \\pm 0.14{\\rm (stat)} \\pm {\\rm 0.48(sys)}$ for the first year sample from the Sloan Digital Sky Survey II (SDSS-II) Supernova Survey. In a recent near-infrared study of nearby SNe,  \\citet{2010AJ....139..120F} found $R_V \\approx 1-2$ for their whole sample, but obtained $R_V \\approx 3.2$ if they exclude their most reddened objects. Alternatively, in a study of 80 nearby SNe Ia, \\citet{2008A&A...487...19N} found a value of $R_V=1.75\\pm0.27$ for their whole sample and a lower value of $R_V\\sim1$ if they restrict the sample to low reddening values. These differences could suggest that the effects of dust may also be dependent on the particular line of sight \\citep{2007ApJ...666..694W} or on the inclination of the host galaxies \\citep{2010arXiv1001.1744M}.  Secondly, the details of the supernova progenitor system could systematically vary between the different galaxy types. Our present theoretical understanding of SNe Ia suggests they are the thermonuclear explosion of a carbon-oxygen white dwarf which has reached the Chandrasekhar limit \\citep{1973ApJ...186.1007W, 2000ARA&A..38..191H}. The mechanism for how the progenitor system accretes mass could be different between galaxy types, either accretion from a nearby companion star (which could have different metallicities depending on the stellar populations in the host galaxy types) or the merger with another white dwarf \\citep{2009ApJ...699.2026R}.  Therefore, there are clear reasons to search for correlations between the properties of Ia's and the properties of their host galaxies. For example, there is a well-established difference between the rates of Ia's in passive and star-forming galaxies, potentially indicating two different paths or timescales for Ia's \\citep{1979AJ.....84..985O, 1990PASP..102.1318V, 2005A&A...433..807M, 2006ApJ...648..868S}. There have also been indications that the host galaxy type correlates with the observed residuals on the SN Hubble diagram, even after standardizing each SN Ia \\citep{2003MNRAS.340.1057S, 2008ApJ...685..752G, 2009arXiv0912.0929K}. For example, \\citet{sullivan2010} recently reported that SNe Ia in massive host galaxies are 0.08 magnitudes brighter than those in lower mass hosts after correction for the light curve shape and color (at a statistical significance of 4$\\sigma$). Such correlations would have important consequences for supernova surveys and could improve the use of SNe Ia as ``standard candles\" \\citep{2009ApJ...699L.139W, sullivan2010}.  We investigate the environmental dependencies of SNe Ia by studying the residuals on the Hubble diagram (around the best fit cosmology) as a function of host galaxy type. We further ask if there are differences in the assumed dust law depending on the type of the host galaxy. In Section 2, we outline the data used in this analysis, which is taken from the full SDSS-II Supernova Survey \\citep{2008AJ....135..338F}. This sample of SNe Ia has several advantages for such environmental studies including high survey efficiency, multi-color ($ugriz$) photometry for all host galaxies and a significant cosmological volume, thus providing a fair sampling of the galaxy distribution. Also, the overall SN rate, as a function of galaxy type, has been measured using this data \\citep{smithm}  In Section 2, and Appendix A, we outline the details of our analysis using two, well-known, public light curve fitting procedures; SALT2 \\citep{2007A&A...466...11G} and MLCS2k2 \\citep{2007ApJ...659..122J, 2009ApJS..185...32K}. We also describe our methodology for defining passive and star-forming host galaxies.  In Section 3, we present our main results, while in Section 4, we discuss these results in light of other work in the field. We conclude in Section 5. ",
        "conclusions": "We present an analysis of the host galaxy dependencies for the SDSS-II Supernova Survey.  We have used 361 SNe Ia (see Table \\ref{tbl_sample}) taken from the full three years of this survey, and then applied several data cuts to ensure we have a clean, well-understood, sample of low redshift SNe ($z<0.21$). We have analysed these data using two well-known light--curve fitting routines (SALT2 and MLCS2k2) to demonstrate that our results are not dependent on the details of the light curve analysis. We summarise below the main conclusions of this work:  \\begin{itemize}  \\item We confirm, to high significance, the strong correlation between host galaxy type and the observed width of the light curve, i.e., quick decline--rate SNe (small $x_1$ values in SALT2), favor passive host galaxies, while bright, slower decline SNe Ia (larger $x_1$ values) favor star-forming galaxies. This has been seen before by several authors. However, we find no correlation between the color of individual SNe Ia and their host galaxy, as illustrated in Figure 2.  \\item We find that SNe Ia are $\\simeq0.1$ magnitudes brighter in passive host galaxies after light curve fitting. This effect is true for both SALT2 and MCLS2k2 analyses. The statistical significance of this difference is between $2$ and $3\\sigma$ dependent upon the details of the fitting methodology and the inclusion of outliers in the color and $x_1$ distributions of these data.  \\item We find evidence for differences in the SN color relationship between passive and star--forming host galaxies. For SALT2, we detect differences in $\\beta$, with passive hosts showing $\\beta\\simeq2.5$ and star--forming hosts prefering $\\beta>3$. For MLCS2k2, we see a similar trend for passive hosts preferring a dust law with $R_V\\simeq1$ and star-forming hosts giving $R_V\\sim2$. The significance of these trends depends on the color range considered, but is greater than $3\\sigma$ for the full SN sample considered herein.  \\item We find that the required intrinsic dispersion for passive galaxy hosts is smaller than that needed for the whole SN sample (and for star--forming hosts), e.g., only $\\sigma_{int}=0.13$ mags is required to obtain a reduced $\\chi^2$ close to unity for passive hosts compared to 0.17 mags for the star-forming sample. This lower intrinsic dispersion for passive hosts is true for both SALT2 and MLCS2k2 light curve fitters.  \\item We demonstrate that the dependence on host galaxy type can be parameterized using the stellar mass of the host galaxies. We show that a 4--parameter fit to the distance modulus of SNe Ia ($M$, $\\beta$, $\\alpha$, $\\gamma$) -- where $\\gamma$ scales with stellar mass -- is better than the usual 3--parameter model given in Eqn. 1. For the data in Figure \\ref{fig9}, we find $\\gamma=0.069\\pm0.014$, or a $4\\sigma$ detection of this parameter. \\end{itemize}  These conclusions are in good agreement with other work, especially Kelly et al. (2009) and Sullivan et al. (2010). In particular, Sullivan et al. (2010) see the same trends in both $M$ and $\\beta$ discussed herein with a similar level of statistical significance. This indicates that these trends are common to several SN surveys and appear to not change significantly with redshift. One possible cause for these correlations is a difference in the host galaxy metallicity (see Gallagher et al. 2008), which is correlated with the host stellar mass and host star-formation activity, and could affect the metallicity of the progenitor star thus leading to changes in the peak brightness of SNe Ia. However, the origin of these correlations requires further study, especially to ensure deficiencies in the light curve fitting techniques are not directly responsible.  The host galaxy dependencies presented in this paper will be important for future supernova cosmology surveys, which may wish to exploit these dependencies to minimize the scatter on the SN Hubble diagram. This could be achieved by including further parameters in the light curve fitting or distance modulus calculation (e.g., Eqn. 2). We will explore these issues in the future with the full SDSS-II SN Survey similar to the recent analysis of Sullivan et al. (2010) for the SNLS."
    },
    "1005/1005.3919_arXiv.txt": {
        "abstract": "{Two main scenarios for the formation of the Galactic bulge are invoked, the first one through gravitational collapse or hierarchical merging of subclumps, the second through secular evolution of the Galactic disc.} {We aim to constrain the formation of the Galactic bulge through studies of the correlation between kinematics and metallicities in Baade's Window ($l=1\\degr$, $b=-4\\degr$) and two other fields along the bulge minor axis ($l=0\\degr$, $b=-6\\degr$ and $b=-12\\degr$).} {We combine the radial velocity and the [Fe/H] measurements obtained with FLAMES/GIRAFFE at the VLT with a spectral resolution of R=20000, plus for the Baade's Window field the OGLE-II proper motions, and compare these with published N-body simulations of the Galactic bulge.} {We confirm the presence of two distinct populations in Baade's Window found in Hill et al. 2010: the metal-rich population presents bar-like kinematics while the metal-poor population shows kinematics corresponding to an old spheroid or a thick disc one. In this context the metallicity gradient along the bulge minor axis observed by Zoccali et al. (2008), visible also in the kinematics, can be related to a varying mix of these two populations as one moves away from the Galactic plane, alleviating the apparent contradiction between the kinematic evidence of a bar and the existence of a metallicity gradient.} {We show evidences that the two main scenarios for the bulge formation co-exist within the Milky Way bulge.} ",
        "introduction": "Although the Milky Way bulge is our closest opportunity to study in detail such a complex chemo-dynamical system, its formation and evolution is still poorly understood. Indeed the high extinction, the crowding, and the superposition of multiple structures along the line of sight make studies of the inner Galactic regions challenging. Two main scenarios have been invoked for bulge formation. The first one is gravitational collapse \\citep{Eggen62} or hierarchical merging of subclumps (\\cite{Noguchi99}, \\cite{Aguerri01}). In this case the bulge formed before the disc and the star-formation time-scale was very short. The resulting stars are old ($>$ 10 Gyr) and have enhancements of $\\alpha$ elements relative to iron which are characteristics of classical bulges. The other scenario is secular evolution of the disc through a bar forming a pseudo-bulge (see, e.g. \\cite{Combes90}, \\cite{Norman96}, \\cite{Kormendy04}, \\cite{Athanassoula05}). After the bar formation  it heats in the vertical direction (\\cite{Combes81}, \\cite{Merrifield96}) giving rise to the typical boxy/peanut aspect. Observational data of individual stars in our Galaxy provides evidence of both scenarios. The presence of a bar in the inner Galaxy has been first suggested by \\cite{Vaucouleurs64} from gas kinematics and confirmed since then by numerous studies including infrared luminosity distribution (e.g. \\cite{BlitzSpergel91}), photometric data (e.g. \\cite{Stanek94}, \\cite{Babusiaux05}, \\cite{Benjamin05}), OH/IR and SiO-maser kinematics (e.g. \\cite{Habing06}), stellar kinematics (e.g. \\cite{Rangwala09a}) and microlensing surveys (e.g. \\cite{Alcock00}). Several triaxial structures may even coexist within the Galactic inner regions (\\cite{Alard01}, \\cite{Nishiyama05}, \\cite{LopezC07}). The boxy aspect of the bulge, detected in near-infrared light profile \\citep{Dwek95}, also argues for a secular formation of the bulge. On the other hand, from medium and high resolution spectroscopic data enhancements of $\\alpha$ elements have been observed (\\cite{McWilliam94}, \\cite{Zoccali06}, \\cite{Fulbright07},  \\cite{Lecureur07}) suggesting a short formation time-scale. The discovery of a possible dwarf remnant in the bulge by \\cite{Ferraro09} also argues for a hierarchical merging scenario. A metallicity gradient has been observed in the minor-axis direction of the bulge \\citep{PaperI} favouring a bulge formation through dissipational collapse. Recently some observed similarities between the [$\\alpha$/Fe] bulge trend and the trend of the local thick stars (\\cite{Melendez08}, \\cite{Ryde09}, \\cite{AlvesBrito10}) suggest that the bulge and the local thick disc experienced similar chemical evolution histories. In what concerns the age, colour-magnitude  diagrams show that most of the Galactic bulge stars are older than 10 Gyr (\\cite{Ortolani95}, \\cite{Feltzing00}, \\cite{Zoccali03}, \\cite{Clarkson08}) although an intermediate-age population exists (\\cite{vanLoon03}, \\cite{GroenewegenBlommaert05}).\\par  Chemo-dynamical modelling of disc galaxy formation in a Cold Dark Matter (CDM) universe stresses the fact that both types of bulges can coexist in the same galaxy. In the \\cite{Samland03}  simulation the bulge contains two stellar populations, an old population formed during the collapse phase and a younger bar population, differing in the [$\\alpha$/Fe] ratio. \\cite{Nakasato03} suggest that the Galactic bulge may consist of two chemically different components: one rapidly formed through subgalactic clump merger in the proto-Galaxy, and the other one formed later in the inner region of the disc. The [Fe/H] abundance of the merger component tends to be smaller than that of the second component. Recently \\cite{Rahimi09}  simulated   bulges formed through multiple mergers and analysed the chemical and dynamical properties of accreted stars with respect to locally formed stars: accreted stars tend to form in early epochs and have lower [Fe/H] and higher [Mg/Fe] ratios, as expected.\\par  Exploring correlations of abundances and kinematics in Baade's Window, \\cite{Soto07}  suggest  a transition in the kinematics of the bulge, from an isotropic oblate spheroid to a bar, at [Fe/H] $= -0.5$ dex. Their  study was based on 315 K and M giants with proper motions of \\cite{Spaenhauer92}, radial velocities of \\cite{Sadler96} and low-resolution abundances of \\cite{Terndrup95} re-calibrated with the  iron abundances scale of \\cite{Fulbright06}. Similar conclusions were obtained  earlier by   \\cite{Zhao94} using a sample of 62 K giants who pointed out the triaxiality of the bulge from kinematics in  Baade's Window and the fact that the metal poor and metal rich populations were not drawn from the same distribution. \\par  In the present paper we analyse the correlations between kinematics and metallicity along  three different minor-axis fields: Baade's Window ($b=-4\\degr$), $b=-6\\degr$ and $b=-12\\degr$,  for which [Fe/H] abundances (\\cite{PaperI} hereafter Paper\\,I and \\cite{PaperII} hereafter Paper\\,II) and radial velocities data for about 700 stars have been determined.  Paper\\,II obtained [Fe/H] and [Mg/H] metallicity distributions for red clump stars in Baade's Window and showed that the sample seems to be separated into two different populations in the metallicity distributions. We here use the kinematic properties of the samples to confirm and constrain the nature of these two populations in Baade's Window and study their relative proportion change along the bulge minor axis fields of Paper\\,I. The paper is organized as follows: section 2 summarizes the data used in this study. In section 3 we analyse Baade's Window by combining our spectroscopic data with OGLE proper motion data. In section 4 we analyse the radial velocity versus metallicity trend along the bulge minor axis. In section 5 we compare our data with some published N-body models. Section 6 discusses our results. ",
        "conclusions": "\\begin{center} \\begin{tabular}{lrrrrcccc} \\hline\\hline Field & $l$  & $b$ & $z_{GC}$ & $N$ & $\\langle V_r \\rangle$ & $\\sigma_{r}$ & skew & kurtosis\\\\ & (\\degr) & (\\degr) & (pc) & & (km/s) &  (km/s) \\\\ \\hline $b = -4\\degr$ & 1.1 & $-4.0$ & 606 & 194 & { }~{ }~5 $\\pm$  7 & 104  $\\pm$  5 &  $-0.07  \\pm  0.1$ &  $-0.54  \\pm  0.2$ \\\\ $b = -6\\degr$ & 0.2 & $-6.0$ & 844 & 201 & $-10 \\pm 6$ & { }~83 $\\pm$ 4 & { }~0.01  $\\pm$  0.1 & $-0.19  \\pm  0.2$ \\\\ $b = -12\\degr$ & 0.0 & $-12.0$ & 1700 & 99 & $-14 \\pm 8$ & { }~80 $\\pm$ 8 & { }~0.38  $\\pm$  0.4 & { }~1.71  $\\pm$  0.7 \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  We first compare the kinematic behaviour along the bulge minor axis (Table \\ref{tab:fieldsSummary}) with previous results from the literature. We saw in section 3 that our Baade's Window radial velocity dispersion is in excellent agreement with previous measurements. At $b=-6\\degr$ our radial velocity dispersion is lower than the BRAVA measure ($\\sigma_{r}=108 \\pm 7$ km/s, \\cite{Howard08}). The decrease of the radial velocity dispersion with galactic latitude we observe is consistent with the SiO maser measurements of \\cite{Izumiura95paperIII}.  All the fields show a negligible skew, while the kurtosis are different. This variation of the kurtosis was not detected in  \\cite{Howard08}. At $b=-12\\degr$ the distribution is significantly pointy (with 99.3\\% confidence according to the Anscombe-Glynn kurtosis test), indicating that the kinematics are significantly affected by the disc. At $b=-6\\degr$ the kurtosis becomes consistent with zero. In Baade's Window the distribution is significantly flattened (with 93\\% confidence in the RGB sample and 99.8\\% in the red-clump sample). \\cite{Sharples90} and \\cite{Rangwala09a} also measured a skew consistent with zero and kurtosis significantly negative in Baade's Window, indicating that the distributions are flat-topped rather than peaked. \\cite{Rangwala09a} concludes that this seems to be consistent with a model of the bar with stars in elongated orbits forming two streams at different mean radial velocities, broadening and flattening the total distribution.  \\begin{figure*} \\begin{minipage}[b]{0.5\\linewidth} \\centering \\includegraphics[width=9cm]{methist.eps} \\caption{Distribution of the metallicity for the different galactic latitudes of Paper\\,I} \\label{fig:metdens} \\end{minipage} \\hspace{0.5cm} \\begin{minipage}[b]{0.5\\linewidth} \\centering \\includegraphics[width=9cm]{metdispVr.eps} \\caption{Dispersion of the radial velocity for the different galactic latitudes as a function of metallicity by bins of 0.4 dex. } \\label{fig:metdispVr} \\end{minipage} \\end{figure*}  We analyse now  kinematic data versus metallicity. The mean radial velocities do not show significant variation with metallicity in any of our fields. Figure \\ref{fig:metdispVr} shows that the velocity dispersion at the rich end decreases significantly with latitude. However the velocity dispersion on the metal-poor end does not vary significantly with latitude. Summing all stars with [Fe/H] $< -0.5$ dex in all fields leads to $\\sigma_{r}=97 \\pm 7$ km/s.  The field $b=-6\\degr$ is close enough to Plaut's Window ($l=0\\degr$, $b=-8\\degr$) so that we can compare our results to the proper motion study of \\cite{Vieira07}. The Besan\\c{c}on model indicates a mean distance of our selected bulge stars to be 6.5 kpc, compatible with the \\cite{Vieira07} selection. Their metallicity distribution is compatible with ours. They do not find neither any change in the proper-motion dispersion as a function of metallicity. They found $\\sigma_l=102\\pm3$ km/s and $\\sigma_b=88\\pm3$ km/s at $b=-8\\degr$, we obtain $\\sigma_r=83$ km/s at $b=-6\\degr$, leading to a full picture coherent with the predictions of the model of \\cite{Zhao96} of $\\sigma_b \\sim \\sigma_r$ and $\\sigma_l > \\sigma_b$ at these latitudes (see the bottom of Fig. 6 of \\cite{Zhao96}).  We have seen in the previous section that Baade's Window kinematic behaviour as a function of metallicity can be interpreted as a mix of two populations, a metal poor component with kinematics that can be associated to an old spheroid population and a metal rich component with bar-driven kinematics. Under this light we can interpret the variation with galactic latitude of both the metallicity distribution function (Fig. \\ref{fig:metdens}) and the kinematic as a function of metallicity (Fig. \\ref{fig:metdispVr}) as the bar population disappearing while moving away from the plane. At high latitudes the foreground disc component dominates the metal rich part of the kinematic behaviour. The metal poor component associated with the old spheroid stays present along the bulge minor axis.  We have obtained an estimation of the variation of the populations with latitude by using the SEMMUL Gaussian components decomposition in the metallicity distribution of the Paper\\,I samples. Table \\ref{tab:b4decomp} gives the decomposition of the Baade's Window RGB sample. The proportion of the metal-rich component is higher than in the red clump sample of Paper\\,II and its spread in metallicity is higher also, most probably due to a difference in the sample selection which is biased towards stars closer to the Sun in the RGB sample. Although both samples are not exactly on the same scale at the high metallicity end (cf section 5.1.4 of Paper II), the mean metallicities for the metal-poor population is extremely similar in both samples (see also the decomposition of the full Baade's Window sample on the same metallicity scale at the end of section 3.1). At $b=-6\\degr$ the Wilks' test allows to keep a solution with 3 components presented in Table \\ref{tab:b6decomp} rather than the 2 components one. Population A and population B could correspond to the population A and B observed in Baade's Window. The mean metallicities are coherent although their spread is smaller. The radial velocity dispersions of population A are identical. The radial velocity dispersion of population B decreases at $b=-6\\degr$ as expected for a bar-like kinematic behaviour (see Fig. \\ref{fig:compFuxSigVr} and associated text). Population C represents only $6\\pm2$\\% of the sample with a low mean metallicity of $-1.1\\pm0.1$ dex and a high velocity dispersion of 127$\\pm$26 km/s, which could therefore be associated to the halo. The Besan\\c{c}on model prediction of only $0.4\\pm0.1$\\% of halo star in the sample could therefore have been underestimated. This population could also have been hidden in the population A at $b=-4\\degr$. Selecting all stars with [Fe/H]$<-0.9$ in our 3 fields we obtain 16 stars with a radial velocity dispersion of $\\sigma_r= 116\\pm21$km/s, which is coherent with the solar neighbourhood velocity dispersions measured in this metallicity range (e.g. \\cite{ChibaBeers00}: $\\sigma_U \\sim 110 \\pm 10$ km/s) containing both thick disc and halo stars. We are not in a position to clearly associate this population either to the halo or to a metal-poor thick disc in our inner galactic samples. SEMMUL did not converge on the $b=-12\\degr$ field due to the smaller number of stars and the higher contamination with thin, thick discs and halo stars expected in this field. At $b=-12\\degr$ the metal rich component present at $b=-4\\degr$ and $b=-6\\degr$ seems to have fully disappeared. The metal rich velocity part of the velocity dispersion corresponds to a disc like component (see next section) while the metal poor part shows a velocity dispersion still coherent with the metal poor population of $b=-4\\degr$ and $b=-6\\degr$, although we cannot distinguish a spheroid and a thick disc contribution.   \\begin{table*} \\caption{SEMMUL Gaussian components decomposition of field $b=-4\\degr$, RGB sample only. The radial velocity dispersion $\\sigma_{r}$ is given in km/s.} \\label{tab:b4decomp} \\begin{center} \\begin{tabular}{lllll} \\hline\\hline Pop & [Fe/H] & $\\sigma_{[Fe/H]}$ & \\% & $\\sigma_{r}$ \\\\ \\hline A  & $-0.31 \\pm 0.05$ & $0.39 \\pm 0.02$ & $33 \\pm 3$ & { }~$88 \\pm 8$ \\\\ B  & { }~$0.13 \\pm 0.02$ & $0.23 \\pm 0.01$ & $67 \\pm 3$ & $109 \\pm 7$ \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  \\begin{table*} \\caption{SEMMUL Gaussian components decomposition of field $b=-6\\degr$. The radial velocity dispersion $\\sigma_{r}$ is given in km/s. } \\label{tab:b6decomp} \\begin{center} \\begin{tabular}{lllll} \\hline\\hline Pop & [Fe/H] & $\\sigma_{[Fe/H]}$ & \\% & $\\sigma_{r}$ \\\\ \\hline C & $-1.09 \\pm 0.07$ & 0.24 $\\pm$ 0.01 & { }~6 $\\pm$ 2 & 127 $\\pm$ 26 \\\\ A & $-0.27 \\pm 0.02$ & 0.24 $\\pm$ 0.01 & 64 $\\pm$ 3 & { }~83 $\\pm$ 5  \\\\ B &  { }~0.14 $\\pm$ 0.02 & 0.13 $\\pm$ 0.01 & 30 $\\pm$ 3 & { }~70 $\\pm$ 6\\\\ \\hline \\end{tabular} \\end{center} \\end{table*}   \\begin{figure*} \\vspace{0.3cm} \\begin{minipage}[b]{0.49\\linewidth} \\centering \\includegraphics[width=8cm]{compSigVr.eps} \\caption{Radial velocity dispersion along the bulge minor axis compared to the BRAVA data (\\cite{Howard08} and \\cite{Howard09}) and the model of \\cite{Zhao96}} \\label{fig:compSigVr} \\end{minipage} \\hspace{0.2cm} \\begin{minipage}[b]{0.49\\linewidth} \\centering \\includegraphics[width=8cm]{compFuxSigVr.eps} \\caption{Radial velocity dispersion along the bulge minor axis compared with the model of \\cite{Fux99}. The metal poor and metal rich population are defined by the 33\\% and 66\\% quantile of the [Fe/H] distribution.} \\label{fig:compFuxSigVr} \\end{minipage} \\end{figure*}    Our analysis of the kinematics as a function of metallicity in Baade's Window shows that our sample can be decomposed in two distinct populations for which we suggest different formation scenarii. The metal poor component does not show any correlation between its velocity components and is therefore consistent with an isotropic rotating population. Paper\\,II showed that this component is enriched in [Mg/Fe]. We interpret this population as an old spheroid with a rapid time-scale formation. The metal rich component shows a vertex deviation consistent with that expected from tags a population with orbits supporting a bar. Paper\\,II showed that this component has a [Mg/Fe] near solar. We interpret this population as a pseudo-bulge formed over a long time scale through disc secular evolution under the action of a bar. This pseudo-bulge is gradually disappearing when moving away from the Galactic plane.  In this context we can give a new consistent interpretation of the metallicity gradient in the bulge. A metallicity gradient is indeed visible in the bulge when observing further away from the plane than Baade's Window (\\cite{Frogel99} and Paper\\,I), while in the inner regions ($\\vert b  \\vert\\leq4\\degr$) no gradient in metallicity has been found (\\cite{Ramirez00} and \\cite{Rich07}). This can be understood if the bar as well as the old spheroid population are both present in the inner regions, leading to a constant metallicity at $\\vert b \\vert\\leq4\\degr$, while the bar influence gradually fades further away from the plane than Baade's Window. At high latitudes only the old spheroid remains: the mean metallicity of the outer bulge measured by \\cite{IbataGilmore95} of [Fe/H]$\\sim-0.3$ dex corresponds very well to our metal poor population. This scenario is also consistent with the distribution of bulge globular clusters of \\cite{Valenti09} who found no evidence for a metallicity gradient but all their clusters with [Fe/H]$>-0.5$ dex are located within ($\\vert b  \\vert\\leq5\\degr$).  In what concerns the age of the two populations, we expect the spheroid component to be old while the pseudo-bulge component may contain both the old stars of the inner disc redistributed by the bar and younger stars whose formation has been triggered by the bar gas flow. \\cite{vanLoon03} found that although the bulk of the bulge population is old, a fraction of the stars are of intermediate age (1 to 7 Gyr). \\cite{GroenewegenBlommaert05} observed Mira stars of ages 1-3 Gyr at all latitudes from $-1.2$ to $-5.8$ in the OGLE-II data. \\cite{Uttenthaler07} found 4 bulge stars with ages lower than 3 Gyr at a latitude of $b=-10\\degr$. \\cite{Bensby09} found 3 microlensed bulge dwarfs with ages lower than 5 Gyr. 87\\% of the variable stars detected by \\cite{KouzumaYamaoka09} are distributed within $\\vert b  \\vert\\leq5\\degr$ and most of them should be large-amplitude and long-period variables such as Mira variables or OH/IR stars. This intermediate age population has been shown to trace the Galactic bar (\\cite{vanLoon03}, \\cite{Izumiura95paperIII}, \\cite{GroenewegenBlommaert05}, \\cite{KouzumaYamaoka09}), although providing a larger bar angle ($\\sim40\\degr$) than studies based on older tracers such as red clump stars ($\\sim20\\degr$, \\cite{Stanek94}, \\cite{Babusiaux05}). This discrepancy in the bar angle could well be explained if the old tracers probed a mix of spheroid and bar structures while the young tracers only probe the bar one (although biases on the longitude area surveyed needs also to be taken into account, see \\cite{Nishiyama05}). If this intermediate age population was associated to a part of the bar component, their presence in the CMDs would decrease while going away from the plane as the main bar component and would therefore be a small fraction of the CMD of \\cite{Zoccali03} at $b=-6\\degr$. \\cite{Clarkson08} obtained a proper motion decontaminated CMD with a well defined old turn-off in an inner field ($l=1\\degr$, $b=-3\\degr$). However we would expect an intermediate age population associated with the bar to be metal rich, which, due to the age-metallicity degeneracy, would imply that this population could be hidden in the CMD of \\cite{Clarkson08} if its contribution is small enough compared to the bulk of the bulge population. The new filter combination proposed by the ACS Bulge Treasury Programme to break the age-metallicity-temperature degeneracy \\citep{Brown09} should provide new lights on this issue.  We note that in Baade's Window, neither the kinematics nor the chemistry allows to distinguish what we call the old spheroid to the thick disc. The mean metallicity of the solar neighbourhood thick disc is however lower (e.g. \\cite{Fuhrmann08} derived [Fe/H] $=-0.6$ and [Mg/H] $=-0.2$) than the mean metallicity of our metal poor population ([Fe/H]$= -0.27$ dex and [Mg/H]$=-0.04$, Paper\\,II). \\cite{Melendez08}, \\cite{Ryde09}, \\cite{Bensby09} and \\cite{AlvesBrito10} observed similarities between the metallicity of the bulge and the metallicity of thick disc stars for metal poor stars. The sample of \\cite{Ryde09} contains 11 of our stars all with less than solar metallicity. Simulations of the formation of thick stellar discs by rapid internal evolution in unstable, gas-rich, clumpy discs \\citep{Bournaud09} show that thick discs and classical bulges form together in a time-scale shorter than 1 Gyr which allow to explain the observed abundance similarities.  The coexistence of classical and pseudo bulge has been observed in external galaxies (\\cite{Prugniel01}, \\cite{Peletier07}, \\cite{Erwin08}) and obtained by N-body simulations (\\cite{Samland03}, \\cite{Athanassoula05}). The chemical and dynamical model of \\cite{Nakasato03} suggests the presence in the bulge of two chemically different components as we found, one formed quickly through the subgalactic clump merger in the proto-Galaxy, and the other has formed gradually in the inner disc. But they do not have a bar in their model. They fitted well the kinematics of \\cite{Minniti96} who observed at $l=8\\degr$, $b=7\\degr$ a decrease of $\\sigma_r$ with metallicity, which is coherent with what we observed at $b=-12\\degr$ and the fact that \\cite{Nakasato03} defined the bulge radius as R$<$2kpc. The chemo-dynamical model of \\cite{Samland03} predicts the different characteristics of our sample: their total bulge population contains two stellar populations: a metal rich population ([Fe/H]$>$0.17) with [$\\alpha$/Fe]$<$0 associated with the bar, and an old population that formed during the proto-galactic collapse, with a high [$\\alpha$/Fe] and a [Fe/H] corresponding to the ``thick disc'' component. Their model also predicts the resulting apparent metallicity gradient along the bulge minor axis (their Fig. 13).  Our study highlights the importance to combine metallicity to 3D-kinematic information to disentangle the different bulge populations. This approach needs to be extended to various galactic longitudes. Gaia will not only provide those but also allow to determine the distances (probing the different structures along the line of sight and removing distance induced biases) and work on an impressively large sample of un-contaminated bulge stars."
    },
    "1005/1005.3416_arXiv.txt": {
        "abstract": "The main objective of this study is to better understand how magnetic helicity injection in an active region (AR) is related to the occurrence and intensity of solar flares. We therefore investigate magnetic helicity injection rate and unsigned magnetic flux, as a reference. In total, 378 ARs are analyzed using $SOHO$/MDI magnetograms. The 24 hr averaged helicity injection rate and unsigned magnetic flux are compared with the flare index and the flare-productive probability in next 24 hr following a measurement. In addition, we study the variation of helicity over a span of several days around the times of the 19 flares above M5.0 which occurred in selected strong flare-productive ARs. The major findings of this study are as follows: (1) for a sub-sample of 91 large ARs with unsigned magnetic fluxes in the range from 3 to 5$\\times$10$^{22}$ Mx, there is a difference in magnetic helicity injection rate between flaring ARs and non-flaring ARs by a factor of 2; (2) the $GOES$ C-flare-productive probability as a function of helicity injection displays a sharp boundary between flare-productive ARs and flare-quiet ones; (3) the history of helicity injection before all the 19 major flares displayed a common characteristic: a significant helicity accumulation of (3--45)$\\times$10$^{42}$ Mx$^2$ during a phase of monotonically increasing helicity over 0.5--2 days. Our results support the notion that helicity injection is important in flares, but it is not effective to use it alone for the purpose of flare forecast. It is necessary to find a way to better characterize the time history of helicity injection as well as its spatial distribution inside ARs. ",
        "introduction": "A solar flare is a sudden, rapid, and intense brightening usually in a magnetically complicated solar active region (AR). Solar flares produce high energy particles, radiation, and erupting magnetic structures that are related to geomagnetic storms. Their strong electromagnetic radiations from radio waves to gamma-rays have direct effect on cell phones and the global positioning system and heat up the terrestrial atmosphere within minutes so that satellites drop into lower orbits (Schwenn 2006). Enormous economic and commercial losses can be caused by these effects (Baker 2004). Therefore there have been significant efforts to develop flare forecasting systems as part of the space weather service. It is generally thought that flare-productive ARs exhibit complex and non-potential magnetic structures related to the stored magnetic energy to power flares. For this reason, many studies of relationship between the solar flare and photospheric magnetic field properties have been carried out since the flare was first observed and recorded by Carrington (1859) and Hodgson (1859). Some examples include the unbalanced changes in the photospheric line-of-sight magnetic field (Cameron \\& Sammis 1999; Spirock et al. 2002; Wang et al. 2002); rapid changes of the sunspot structure associated with a substantial fraction of flares (Liu et al. 2005; Deng et al. 2005; Wang et al. 2004a, 2005; Chen et al. 2007); the magnetic shear angle evolution (Hagyard et al. 1984; Hagyard \\& Rabin 1986; Sivaraman et al. 1992; Schmieder et al. 1994; Wang et al. 1994, 2004a); the horizontal gradient of longitudinal magnetic fields (Zirin \\& Wang 1993; Zhang et al. 1994; Tian et al. 2002); electric current (Canfield et al. 1993; Lin et al. 1993); magnetic helicity injection (Moon et al. 2002a, 2000b; Sakurai \\& Hagino 2003; Yokoyama et al. 2003; Park et al. 2008). Based on the above mentioned studies, current flare forecasting models are moving toward multiple-magnetic parameter-based approaches (Leka \\& Barnes 2003a, 2003b; Li et al 2008) from sunspot-morphological evolution-based approaches (McIntosh 1990; Gallagher et al. 2002).  Magnetic helicity is a measure of twists, kinks, and inter-linkages of magnetic field lines (Berger \\& Field 1984; Pevtsov 2008) and is a useful and important parameter to indicate topology and non-potentiality of a magnetic field system. Naturally magnetic helicity studies have been conducted to the energy buildup and instability leading to solar eruptions and coronal mass ejections (e.g., Rust 2001; Kusano et al. 2004; Phillips et al. 2005; Fan 2009). There were a number of studies related to a rapid magnetic helicity change as an impending condition or a trigger for solar flares (e.g. Moon et al. 2002a, 2002b). LaBonte et al. (2007) surveyed magnetic helicity injection in 48 X-flaring ARs and 345 non-X-class flaring regions, and found that a necessary condition for the occurrence of an X-class flare is that the peak helicity flux has a magnitude $\\textgreater$ 6$\\times$10$^{36}$ Mx$^2$ s$^{-1}$. Park et al. (2008) found that a substantial amount of helicity is accumulated before the flare in all the 11 X-class flare events, and suggested a warning sign of flares can be given by the presence of a phase of monotonically increasing helicity. Motivated by these results we explore the feasibility of using magnetic helicity to build a flare forecasting system utilizing a data sample covering almost one solar cycle. ",
        "conclusions": "We have investigated the time variations of $\\dot{H}$ and $\\Phi$ in 378 solar ARs and compared the two average parameters, $|$$<$$\\dot{H}$$>$$|$ and $<$$\\Phi$$>$, with $F_{idx}$. Although there is a large amount of scatter in the data samples, we found a moderate correlation between the parameters and $F_{idx}$. The larger $F_{idx}$ an AR has, the larger values of $\\dot{H}$ and $\\Phi$ it presents. To improve the correlation, we have defined a new parameter as an equally weighted linear combination of the two rescaled parameters (0.5 of each). The logarithmic-scale CC of $F_{idx}$ versus the new parameter increased slightly to 0.47. It is not surprising because $|$$<$$\\dot{H}$$>$$|$ is well correlated with $<$$\\Phi$$>$ as shown in Figure 3. Moreover, by considering 48 and 72 hr profiles of $\\dot{H}$ and $\\Phi$ for calculation of the two average parameters, we have executed the same correlation study between the two parameters and the next-day flare index. We found that the longer the period we use for average, the worse the correlation will be, especially in the case of $|$$<$$\\dot{H}$$>$$|$ (CC=0.38 and 0.21 for 48 and 76 hr periods, respectively). This might be because we do not consider the flaring history before or during the measurement period of the parameters, but we compare the parameters with $F_{idx}$ calculated only for the following day of the measurement period.  We understand that no matter which method is used, the correlation between the parameters and flare index is not high. This is an intrinsic problem for flare forecasting as the occurrence of a flare depends not only on the amount of magnetic energy built up in an AR, but also on how it is triggered. For example, if new flux-rope emergence is the driver of flares (Schrijver 2009), or if a flare is exactly a result of a small and localized (quite possibly unobservable) perturbation affecting the whole system like self-organized criticality dynamics (Bak et al. 1987; B\\'{e}langer et al. 2007), then it is not feasible to carry out prediction of flare onset time and magnitude by using present-day parameters derived from photospheric magnetic field observations. More specifically for our case, helicity accumulation might be a necessary, but not sufficient condition for flares. Perhaps a triggering mechanism is necessary even a magnetic system has enough non-potentiality to power a flare (so-called metastable state). This idea agrees with the study that a number of X-class flares occurred during the phase of almost constant helicity after the phase of 2--3 days of monotonically increasing helicity (Park et al. 2008).  Interestingly, contrary to the expectation that magnetic helicity injection is more closely related to flare productivity than to magnetic flux, our result shows that the correlation between $|$$<$$\\dot{H}$$>$$|$ and $F_{idx}$ is not stronger than that between $<$$\\Phi$$>$ and $F_{idx}$. The logarithmic-scale CCs of $F_{idx}$ versus $|$$<$$\\dot{H}$$>$$|$ and $<$$\\Phi$$>$ are 0.42 and 0.43, respectively. If only the flaring groups with non-zero flare index are considered, then $|$$<$$\\dot{H}$$>$$|$ is not better than $<$$\\Phi$$>$ in predicting how strong the flares will be. This might be due to the fact that we simply use the 1 day average of $\\dot{H}$ in the entire AR for comparison with $F_{idx}$ without more specifically characterizing the temporal and spatial evolution of helicity in the AR related to a flaring condition. Magnetic helicity, however, is useful in predicting whether an AR will produce flares or not. Note that predicting the occurrence of flares is different from predicting the strength of flares. By examining more careful studies such as the helicity injection difference between flare-productive and flare-quiet ARs, the flare-productive probability as a function of $|$$<$$\\dot{H}$$>$$|$, and the temporal evolution of helicity in major flare-producing ARs, we have found that magnetic helicity injection has some interesting features related to flares as follows.  \\begin{enumerate} \\item For 91 AR samples in the range (3--5)$\\times$10$^{22}$ Mx of large $<$$\\Phi$$>$ the $flaring$ group has $|$$<$$\\dot{H}$$>$$|$ about twice greater than that of the $non$-$flaring$ group. On the other hand, 118 AR samples of large $|$$<$$\\dot{H}$$>$$|$ do not show the significant difference in $<$$\\Phi$$>$ between the $flaring$ and $non$-$flaring$ groups. \\item The helicity parameter $|$$<$$\\dot{H}$$>$$|$ demonstrates a rapid increase of $P_{i}$ compared to that of $<$$\\Phi$$>$ in the rescaled range of 0--0.15 of the parameter. $P_{C}$($|$$<$$\\dot{H}$$>$$|$), especially, quickly reaches up to $\\sim$90$\\%$ from 46$\\%$ in the very low rescaled range, 0--0.15, of the parameter, and it retains a high value above 90$\\%$ in the rest, 0.15--0.6, of the statistically meaningful range. \\item Helicity of (3--45)$\\times$10$^{42}$ Mx$^2$ accumulates significantly and consistently over 0.5--2 days for all the 19 major flares under investigation. More specifically, following the significant amount of long-term helicity accumulation with fast injection rate, 4 and 8 flares occurred when helicity injection rate starts to become slow (sometimes almost zero) and reverse its sign, respectively. \\end{enumerate}  Based on these results, the magnetic helicity can be used for the improvement of flare forecasting. First of all, when an AR has large $<$$\\Phi$$>$, we may better determine whether or not it will produce a flare by considering $|$$<$$\\dot{H}$$>$$|$ of the AR. Second, the helicity parameter $|$$<$$\\dot{H}$$>$$|$ would allow us to establish a better defined cutoff between C-flare-productive and C-flare-quiet ARs than $<$$\\Phi$$>$ if we take into account a sharp increase of $P_{i}$ in the very low rescaled range of the parameter. Third, an early warning sign of flare occurrence could be based on tracking of a phase of monotonically increasing helicity because there is always a significant amount of helicity accumulation a few days before major flares. We might also make an urgent warning sign when helicity injection rate becomes very slow or the opposite sign of helicity starts to be injected after the significant helicity accumulation phase. The sign reversal of the magnetic helicity may support the numerical simulation model for solar flare onset proposed by Kusano et al. (2003b) in which they showed that magnetic reconnection quickly grows in the site of the helicity annihilation with different signs. Some observations of helicity inversion, similar to our result, were also reported around the time of flare onset (Kusano et al. 2003a; Yokoyama et al. 2003; Wang et al. 2004b). For more practical and advanced flare forecasting, we may think how to consider the past history of flare occurrence in an AR under investigation and combine the helicity parameter with others with different weighting coefficients. Besides that it would be required to better characterize not only the time history of helicity injection but also its spatial distribution inside ARs.  Finally, our study may lead to some physical understanding of flare on-set. For example, why do only some of the samples with the large helicity injection produce major flares, but not for all? Is a significant amount of helicity accumulation necessary or sufficient conditions for flares? We believe that the study of magnetic helicity in a coronal volume of an AR will help us to better explain physically for these questions and understand pre-flare conditions and energy storage process of flares in more detail. Our future work to calculate the coronal helicity by using three-dimensional non-linear force-free magnetic field extrapolations is therefore in progress."
    },
    "1005/1005.4693_arXiv.txt": {
        "abstract": "We recently proposed that molecular cloud dense cores undergo a prolonged period of quasi-static contraction prior to true collapse. This theory could explain the observation that many starless cores exhibit, through their spectral line profiles, signs of inward motion. We now use our model, together with a publicly available radiative transfer code, to determine the emission from three commonly used species - N$_2$H$^+$, CS, and HCN. A representative dense core of $3\\,\\,\\Msun$ that has been contracting for 1~Myr has line profiles that qualitatively match the observed ones. In particular, optically thick lines have about the right degree of blue-red asymmetry, the empirical hallmark of contraction. The \\hbox{$J\\,=\\,2\\rightarrow\\,1$} rotational transition of CS only attains the correct type of profile if the species is centrally depleted, as has been suggested by previous studies. These results support the idea that a slow, but accelerating, contraction leads to protostellar collapse. In the future, the kind of analysis presented here can be used to assign ages to individual starless cores. ",
        "introduction": "Over the last decade, astronomers have carefully scrutinized the large fraction of molecular cloud dense cores that contain no point sources of radiation, in order to understand more fully the onset of star formation. It was already established that the gross properties of these objects differ little from their counterparts with internal stars \\citep{lm99}, except perhaps for a mass that is slightly higher on average \\citep[][Fig.~3]{jm99}. More detailed spectroscopic studies of molecular line emission revealed that many starless cores exhibit signs of inward contraction \\citep{wi99,lmt99,ge00, lmt01,sc07}. The hallmark of this so-called infall signature is an asymmetric emission line, often with a self-absorption dip, that is skewed to the blue.  There is a variety of line profiles seen in dense cores, sometimes even in the same object, and the profile just described is by no means universal. It is preferentially found in starless cores with a higher central column density and optical depth \\citep[e.g.][]{lmt99,ge00}. Such cores also exhibit chemical abundances that differentiate them from lower-density objects \\citep{c05,kc08}. The kinematic signature is especially common in more evolved dense cores that contain Class~0 and Class~I sources, i.e., very young, embedded stars \\citep{me00}. These facts, together with the simple and compelling infall interpretation of the asymmetric profiles \\citep{m96}, have convinced most researchers that we are witnessing a type of contraction that precedes the free-fall collapse onto a protostar. The contraction hypothesis is especially convincing for starless cores, where there is no possibility of confusion from stellar outflows.  What is the physical cause of this motion? The observed magnitude of the infall velocity is typically a few tenths of the sound speed. This finding rules out ambipolar diffusion as the underlying process, since the characteristic ion-neutral drift speed is much smaller \\citep[e.g.][]{cb01}. Conversely, the subsonic level of the velocity, together with the statistical prevalence of starless cores \\citep[e.g.][]{f09}, indicate that such objects are relatively long-lived \\citep{ki05}, and thus are unlikely to be in a true state of collapse, as proposed by \\citet{kf05}.  In a recent contribution \\citep[][hereafter Paper~I]{sy09} we introduced a new model for the contraction. Starless cores, like all self-gravitating objects, undergo bulk oscillations. Indeed, there are now several cases where spectroscopic observations indicate that such oscillations are occurring \\citep{r06,b07,v08}. For an object that has evolved to the point of collapse, by whatever means, the frequency of the lowest normal mode vanishes. We studied, using perturbation theory, how a cloud subject to such a frozen (more commonly called `neutral') mode slowly evolves toward full-blown collapse. We found that this previously unrecognized phase of quasi-static contraction lasts a considerable period (of order $10^6$~yr), during which time the cloud attains internal speeds similar to those observed (see Fig.~4 of Paper~I).  Encouraged by this result, we now seek to forge a stronger link to the observations. We calculate, using a publicly available radiative transfer code, the emitted line spectrum from our model cloud. We choose three molecules that are among the most commonly employed in the spectroscopic studies: N$_2$H$^+$, CS, and HCN. Following the early surveys of \\citet{tt77} and \\citet{w92}, N$_2$H$^+$ is often used as a tracer of dense, relatively quiescent cloud gas. Its rotational lines are optically thin in low-mass dense cores. Thus, they do not exhibit either self-absorption or blue-red asymmetry, but serve principally as a gauge of excitation temperature and as a comparison benchmark with lines that are sensitive to internal cloud motion. The rotational lines of CS are often optically thick, and can have asymmetric profiles that have long been interpreted kinematically \\citep{w86}. On the other hand, the infall signature is relatively weak in starless cores, in part because the polar molecule CS is depleted at high density from the gas phase onto grain mantles \\citep{bl97}. Finally, the \\hbox{$J\\,=\\,1\\rightarrow\\,0$} rotational transition of HCN has the strongest infall signature of all, and the molecule appears to suffer little depletion \\citep{a05,s07}.  Although our model cloud is a highly idealized isothermal sphere, the results of this study further lend credence to the underlying dynamical picture. For a representative 3\\,\\,$\\Msun$ cloud that has been contracting for 1~Myr, we find that all the calculated line profiles resemble those observed, at least in an average sense. Interestingly, the CS profile is only acceptable once we include central depletion of the molecule. In Section~2 below, we first describe our implementation of the radiative transfer program. Section~3 then summarizes observations to date of the three molecules in starless dense cores. We describe, in Section~4, how the emergent line profiles alter when we change our physical input parameters. Judging from these trends, we then select a canonical model cloud and compare its emission spectrum to the available data. Finally, Section~5 discusses limitations of the dynamical theory and the prospects for improvement. We also indicate how the model, even in its current, simplified form, can be used to assign contraction ages to individual starless cores. ",
        "conclusions": "This suite of radiative transfer calculations performed on our theoretical cloud model has yielded two important findings. First, we have been able to match the generic, observed properties of the selected line profiles using a reasonable cloud mass, turbulent velocity dispersion, and evolutionary time. This success may indicate that the model has captured the essential physics of the contraction process itself. Note also that the accelerating character of the contraction in our theory explains naturally the key observation that starless cores of lower density exhibit less asymmetry in their line profiles \\citep{ge00,wi06}. Second, the sensitivity of the emergent line profiles to the input parameters means that the model offers a practical route to gauge starless core properties that have previously been difficult to ascertain. For example, given the empirical mass and kinetic temperature of an object, the matching of line profiles would yield a first estimate of its contraction age. Obtaining ages for numerous starless cores in different environments would aid greatly in understanding the onset of star formation.  Of course, the observations offer more detail than the model can explain, at least in its current, highly idealized form. We cannot yet account for the rather surprising spatial extent of asymmetric CS emission \\citep{me00,lmt01}. Nor can we explain the relative intensities of the HCN hyperfine lines \\citep{s04}. Explaining both features of the data may require alterations in the spatial abundance gradients, in the manner employed in Sections 4.2 and 4.4. Additionally, an improved model would consider more carefully how the interior, relatively quiescent gas joins onto surrounding, more turbulent, material. In fact, it has long been realized that the pattern of HCN hyperfine anomalies requires the presence of an envelope of relatively high excitation temperature \\citep{g93}. \\citet{go09} have recently simulated the birth of dense cores out of spherical, converging flows. Such calculations, when extended to three dimensions, may offer additional insight into the nature of the core-envelope transition.  Finally, we must face the obvious fact that real starless cores are not spherical, as the model currently assumes. The shapes of these objects may be partially determined by the anisotropic pressure of their turbulent environments, as found in recent simulations \\citep[e.g.][]{o08}. Another important factor setting the gross core properties is the force associated with the internal magnetic field \\citep{l04}. It would be instructive to study, using the method of Paper~I, the onset of collapse in marginally stable, magnetostatic equilibria, such as those found by \\citet{t88}. Such an investigation would add a significant degree of realism to our picture of protostellar collapse."
    },
    "1005/1005.3599_arXiv.txt": {
        "abstract": "{The exact period determination of a multi-periodic variable star based on its luminosity time series data is believed a task requiring skill and experience. Thus the majority of available time series analysis techniques require human intervention to some extent.} {The present work is dedicated to establish an automated method of period (or frequency) determination from the time series database of variable stars.} {Relying on the \\sc SigSpec \\rm method (Reegen 2007), the technique established here employs a statistically unbiased treatment of frequency-domain noise and avoids spurious (i.\\,e. noise induced) and alias peaks to the highest possible extent. Several add-ons were incorporated to tailor \\sc SigSpec \\rm to our requirements. We present tests on 386 stars taken from ASAS2 project database.} {From the output file produced by \\sc SigSpec, \\rm the frequency with maximum spectral significance is chosen as the genuine frequency. Out of 386 variable stars available in the ASAS2 database, our results contain 243 periods recovered exactly and also 88 half periods, 42 different periods etc.} {\\sc SigSpec \\rm has the potential to be effectively used for fully automated period detection from variable stars' time series database. The exact detection of periods helps us to identify the type of variability and classify the variable stars, which provides a crucial information on the physical processes effective in stellar atmospheres.} \\keywords { Astronomical instrumentation, methods and techniques -- Stars: variables: general } ",
        "introduction": "Iterative measurements of a quantity over time yield a time series. This work deals with time series obtained from photometric observations of variable stars, which is the photometric flux (or magnitude) versus time, with or without data gaps. The majority of astronomical measurements cannot be taken continuously over long periods of time due to several reasons. For ground based observations, daylight and the weather conditions are unavoidable sources of data gaps, and observations from space may suffer from cosmic particle impacts or stray light corruption occasionally producing data points beyond repair (Reegen et al. 2006).  The method to extract physical information from time-resolved data is commonly known as time series analysis. The present paper is primarily dedicated to regular variables, which exhibit a strictly periodic photometric signal, so the primary goal of our analysis will be to identify the appropriate frequency in case of a mono-periodic star. If the detected frequency is exact, the intrinsic scatter of a phased light curve will attain a minimum. For multi-periodic variables, a step-by-step prewhitening procedure has established and is widely used in combination with multi-sine least-square fitting techniques (Sperl 1998; Lenz \\& Breger 2005). The identified periods permit to identify the type of variability and deduce important astrophysical parameters. The Fourth Variable Star Working Group meeting held at Geneva Observatory, Switzerland in 2005 elaborates about the Period Search Benchmarks (Laurent Eyer)\\footnote{\\tt http://obswww.unige.ch/\\textasciitilde eyer/VSWG/Meeting4/MINUTES/g\\\\aia-vswg4.html}. ",
        "conclusions": "We investigated several existing programs and software packages which are used for the time series analysis of variable stars and found that human intervention is required in some way to confirm the true period. Full automation could be achieved only with {\\sc SigSpec} combined with few additional subroutines. The majority of results are found to be satisfying with the published frequencies. Thus {\\sc SigSpec} turned out to be a useful utility for future massive variability surveys and also for the re-analysis of already existing time series data of periodic variables."
    },
    "1005/1005.1887_arXiv.txt": {
        "abstract": "{Stellar density and bar strength should affect the temperatures of the cool (T~$\\sim 20-30$~K) dust component in the inner regions of galaxies, which implies that the ratio of temperatures in the circumnuclear regions to the disk should depend on Hubble type.  We investigate the differences between cool dust temperatures in the central 3~kpc and disk of 13 nearby galaxies by fitting models to measurements between 70 and 500~$\\mu$m.  We attempt to quantify temperature trends in nearby disk galaxies, with archival data from \\emph{Spitzer}/MIPS and new observations with \\emph{Herschel}/SPIRE, which were acquired during the first phases of the \\emph{Herschel} observations for the KINGFISH (key insights in nearby galaxies: a far-infrared survey with \\emph{Herschel}) sample.  We fit single-temperature modified blackbodies to far-infrared and submillimeter measurements of the central and disk regions of galaxies to determine the temperature of the component(s) emitting at those wavelengths.  We present the ratio of central-region-to-disk-temperatures of the cool dust component of 13 nearby galaxies as a function of morphological type.  We find a significant temperature gradient in the cool dust component in all galaxies, with a mean center-to-disk temperature ratio of $1.15 \\pm 0.03$.  The cool dust temperatures in the central $\\sim3$~kpc of nearby galaxies are 23 ($\\pm3$)\\% hotter for morphological types earlier than Sc, and only 9 ($\\pm3$)\\% hotter for later types.  The temperature ratio is also correlated with bar strength, with only strongly barred galaxies having a ratio over 1.2.  The strong radiation field in the high stellar density of a galactic bulge tends to heat the cool dust component to higher temperatures, at least in early-type spirals with relatively large bulges, especially when paired with a strong bar.} ",
        "introduction": "The infrared emission from galaxies contains roughly half of the entire energy budget in the Universe (e.g., Hauser \\& Dwek \\cite{hauser01}). In addition to providing information on the amount of attenuation suffered by the stellar light, the infrared emission provides clues to important physical quantities, such as the metal, dust, and cold gas content of galaxies (e.g., Draine et al.\\ \\cite{draine07a}, Bernard et al.\\ \\cite{bernard08}).  Infrared spectral energy distributions (SEDs), especially those extending into the submillimeter regime, can be used to measure the dust mass and temperature in galaxies (e.g., Dunne et al. \\cite{dunne00,dunne01}; Seaquist et al.\\ \\cite{seaquist04}; Vlahakis et al.\\ \\cite{vlahakis05}; Draine et al.\\ \\cite{draine07a}; Willmer et al.\\ \\cite{willmer09}; Liu et al.\\ \\cite{liu10}).  These temperature measurements have shown that the dust is warmer in the centers of galaxies than in the outskirts (e.g., Alton et al.\\ \\cite{alton98}; Radovich et al.\\ \\cite{radovich01}; Melo et al.\\ \\cite{melo02}; Dupac et al.\\ \\cite{dupac03}).  Cool dust at roughly the same temperature in spiral disks is detected {\\it globally} at longer wavelengths (850~$\\mu$m and 1.3~mm; Siebenmorgen et al.\\ \\cite{siebenmorgen99}; Dunne et al.\\ \\cite{dunne00,dunne01}; Vlahakis et al.\\ \\cite{vlahakis05}), but there is also some evidence of a warmer temperature component associated with the central regions.  Warmer dust temperatures tend to be associated with significant star-formation activity, the resulting intense interstellar radiation field (Stevens et al.\\ \\cite{stevens05}) and the earlier Hubble type (Bendo et al.\\ \\cite{bendo03}).  Here we use the infrared SEDs, from 70 to 500~$\\mu$m (a range which should be dominated by emission from grains in thermal equilibrium with the radiation field, e.g., Popescu et al.\\ \\cite{popescu00}, Engelbracht et al.\\ \\cite{engelbracht08}), of a local sample of 13 galaxies spanning the range of spiral galaxies in the Hubble sequence, to derive the temperature of the cool (T$ \\sim 20-30$~K) dust component of the central region and the disks separately, and investigate differences in the dust heating in the two regions.  To achieve this goal, we use the 250, 350, and 500~$\\mu$m \\emph{Herschel}/SPIRE (Spectral and Photometric Imaging REceiver; Griffin et al.\\ \\cite{griffin10}) images of the galaxies combined with the \\emph{Spitzer}/MIPS (Multiband Imaging Photometer for \\emph{Spitzer}; Rieke et~al.\\ \\cite{rieke04}) 70 and 160~$\\mu$m images.  Eventually we will use PACS (Photodetector Array Camera and Spectrometer; Poglitsch et al.\\ \\cite{poglitsch10}) imaging for this study, but at the time of this writing, the data were not yet available.  These galaxies have little nuclear activity that might heat the dust, with only NGC~1097 having a strongly active nucleus, so this is the first study to cleanly separate the far-infrared properties of central and disk regions in a sample of normal galaxies.  Until recently, little work has been done to dissect the dust emission in galaxies into sub-galactic components, owing to the general paucity of infrared images with the required angular resolution and to poor long-wavelength sensitivity.  Some recent work includes a study of the galaxy pair NGC1512/1510 (a target also discussed in this paper), which finds that the dust temperature in the central region of NGC1512 is slightly higher than in the disk and that there is a significantly higher fraction of warm dust, in agreement with the center of NGC1512 being a starburst (Liu et al.\\ \\cite{liu10}).  Work by Mu\\~{n}oz-Mateos et al.\\ (\\cite{munoz-mateos09}) examines radial trends in dust properties in a number of nearby galaxies.  The \\emph{Herschel} Space Telescope promises to yield a breakthrough in the study of subgalactic components in galaxies.  This paper is the first investigation to leverage the longest wavelength \\emph{Herschel} data available for the KINGFISH (key insights on nearby galaxies: a far-infrared survey with \\emph{Herschel}; this program is largely derived from SINGS, the \\emph{Spitzer} infrared nearby galaxy survey by Kennicutt et al.\\ \\cite{kennicutt03}) sample of nearby galaxies, which will eventually total 61. Companion papers from the science demonstration phase for this program showcase the shorter wavelength imaging (Sandstrom et al.\\ \\cite{sandstrom10}) and the spectroscopic data (Beir\\~{a}o et al.\\ \\cite{beirao}).  Here we present new SPIRE images acquired during the first few months of \\emph{Herschel} operations, in the context of the KINGFISH Open-Time Key Project.  We divide each galaxy into two spatially-resolved zones: the circumnuclear region and the surrounding disk.  Then we compare central temperatures with those for the disk. ",
        "conclusions": "We used far-infrared data from \\emph{Spitzer} and submillimeter data from \\emph{Herschel} to compute separate SEDs for the center and disk regions of 13 nearby galaxies observed as part of the KINGFISH program.  We fit those SEDs (at wavelengths longer than 70~$\\mu$m) with blackbody functions (modified by a frequency-dependent emissivity) to compute temperatures.  On average, the cool dust temperature of the central component is $15 \\pm 3$\\% hotter than the disk.  We find that the central temperature is higher than the disk by 20\\% to 50\\% in galaxies of type S0 to Sb, but only 9\\% higher in later types.  This ratio is also higher (at $1.29 \\pm 0.04$) in strongly barred galaxies than in weakly barred galaxies (at $1.09 \\pm 0.03$).  The data therefore indicate that the large (or ``classical'') grains that dominate the far-infrared and submillimeter emission are warmer in the centers of those galaxies with a substantial bulge and/or a strong bar.  This may simply be caused by the higher density of the radiation field in the centers of early-type spirals, enhanced star formation due to the bar, or some combination of the two.  A cleaner separation of morphological components (perhaps with larger samples and/or less distant galaxies) and a more thorough assessment of the density of starlight and star formation activity, plus an evaluation of the impact of central nonthermal sources, may help separate these effects.  The analysis presented here illustrates the power of {\\it Herschel} observations in characterizating the spatially resolved distribution of dust in nearby galaxies.  This power will grow with the use of better-resolved far-infrared SEDs as measured by PACS (Poglitsch et al.\\ \\cite{poglitsch10}), which will let us measure smaller and/or more distant galaxies and determine radial trends of dust temperature."
    },
    "1005/1005.0848_arXiv.txt": {
        "abstract": "{High precision astrometry requires an accurate geometric distortion solution. In this work, we present an average correction for the Blue Camera of the Large Binocular Telescope which enables a relative astrometric precision of $\\sim$15 mas for the $B_{\\rm Bessel}$ and $V_{\\rm Bessel}$ broad-band filters.  The result of this effort is used in two companion papers:\\ the first to measure the absolute proper motion of the open cluster M~67 with respect to the background galaxies;\\ the second to decontaminate the color-magnitude diagram of M~67 from field objects, enabling the study of the end of its white dwarf cooling sequence.  Many other applications might find this distortion correction useful.} ",
        "introduction": "\\label{sec1}  Modern wide field imagers (WFI) equipped with CCD detectors began their operations at the end of the last century, however -- after more than 10 years -- their astrometric potential still remains somehow unexploited (see Anderson et al.\\ \\cite{anderson06}, hereafter Paper~I).  It is particularly timely to begin exploring their full potential now that WFI start to appear also at the focus of the largest available 8m-class telescopes.   The present work goes in this direction, presenting a correction for the geometric distortion (GD) of the Blue prime-focus Large Binocular Camera (LBC), at the Large Binocular Telescope (LBT).  Unlike in Paper~I, in which we corrected the GD of the WFI at the focus of the 2.2m MPI/ESO telescope (WFI@2.2m) with a look-up table of corrections, for the LBC@LBT we will adopt the same technique described in Anderson \\& King (\\cite{AK03}, hereafter AK03), and successfully applied to the new Wide field Camera 3/UV-Optical channel on board the {\\it Hubble Space Telescope} (Bellini \\& Bedin \\cite{bb09}, hereafter BB09).  This article is organized as follows: Section 2 briefly describes the telescope/camera set up; Section 3 presents the data set used.  In section 4, we describe the steps which allowed us to obtain a solution of the GD, for each detector separately, while in Section 5 we presents a (less accurate) inter-chip solution.  Distortion stability is analyzed in Section 6, and a final Section summarizes our results.    \\begin{figure*}[ht!] \\centering \\includegraphics[height=10.2cm]{13783f1a.ps} \\includegraphics[height=10.0cm]{13783f1b.ps}\\\\ \\caption{(\\textit{Left}): LBC-Blue mosaic layout; ``$\\ast$'' marks the center for each chip (see Sect.\\ \\ref{sec:3} for the operative definition of centers and chips), while ``+'' marks the here defined center of the mosaic.  (\\textit{Right}): Each MEF file consists of 4 images. Each image is composed by one scientific region and two overscan regions, covering the first 50 and the last 206 pixel columns (shaded).} \\label{fig1} \\end{figure*} ",
        "conclusions": "\\label{sec:10}  By using a large number of well dithered exposures we have found a set of third-order-correction coefficients for the geometric distortion solution of each chip of the LBC-Blue, at the prime focus of the LBT.  The use of these corrections removes the distortion over the entire area of each chip to an accuracy of $\\sim$0.09 pixel (i.e.\\ $\\sim$20 mas), the largest systematics being located in the 200-400 pixels closest to the boundaries of the detectors.  Therefore, we advise the use of the inner parts of the detectors for high-precision astrometry. The limitation that has prevented us from removing the distortion at even higher level of accuracies -- in addition to atmospheric effects and to the relatively sparsity of the studied field -- is the dependency of the distortion on the scale changes that result from thermal and/or gravitational induced variations of the telescope+optical structure.  If a dozen (or more) well distributed high S/N stars are available within the same chip, a general 6-parameter linear transformation could register relative positions in different images down to about 15 mas.  If the field is even more densely populated, then a local transformation approach [as the one adopted in Bedin et al.\\ (\\cite{bedin03}), from space, or in Paper~I, II, III) from ground] can further reduce these precisions to the mas level.  [Indeed, using these techniques and this very same data-set we were able to reach a final precision of $\\sim$1 mas$\\,$yr$^{-1}$ (Bellini et al., submitted to \\aap\\ Letters)].  These are the precisions and accuracies with which we can hope to bring one image into another image by adopting:\\ {\\it conformal}, {\\it general}, or {\\it local} transformations.  In the case of absolute astrometry, however, the accuracies are much lower.  During the available limited number of nights of observations (and atmospheric conditions), we observed scale-variations up to 5 parts in $10\\,000$, even during the same night.  This implies that astrometric accuracy -- which completely relies on our GD solution -- can not be better than $\\sim$250 mas ($\\sim$1 pixel) within a given chip (from center to corners), and can be as large as $0\\farcs5$ ($\\sim$2 pixels) in the meta-chip system.  This value is in-line with the meta-chip stability observed in other ground-based WFI (Paper~I), and absolutely excellent for a ground-based prime-focus instrument with such a small focal ratio and large FoV.  Thankfully, several stars from astrometric catalogs such as the UCAC-2, GSC-2, 2MASS, will be always available within any given LBC-Blue large FoV. These stars, in addition to provide a link to absolute astrometry (as done for example in Rovilos et al.\\ 2009), will enable constrains of linear terms in our GD solution, and to potentially reach an absolute astrometric precision of 20 mas. The fact that we are able to reach good astrometric precision also for saturated stars will make the comparison between these catalogs and the sources measured in the -- generally deeper -- LBC images, even easier.  For the future, more data and a longer time-baseline are needed to better characterize the GD stability of LBC@LBT detectors on the medium and long time term. This could make it possible to:\\ (1) determine a multi-layer model of the distortion which would properly disentangle the contributions given by optical field-angle distortion, light-path deviations caused by filters and windows, non-flat CCDs, CCDs artifacts, alignment errors of the CCD on the focal plane, etc.; and (2) allow for time-dependent and/or mis-alignments of mirrors, filters/windows, and CCDs."
    },
    "1005/1005.5695_arXiv.txt": {
        "abstract": "The photon polarization operator in superstrong magnetic fields induces the dynamical photon ``mass'' which leads to screening of Coulomb potential at small distances $z\\ll 1/m$, $m$ is the mass of an electron. We demonstrate that this behaviour is qualitatively different from the case of $D=2$ QED, where the same formula for a polarization operator leads to screening at large distances as well. Because of screening the ground state energy of the hydrogen atom at the magnetic fields $B \\gg m^2/e^3$ has the finite value $E_0 = -me^4/2 \\ln^2(1/e^6)$. ",
        "introduction": "The Larmour radius of the electron orbit $a_H = 1/\\sqrt{eB}$ is much smaller than Bohr atomic radius $a_B = 1/(me^2)$ for homogenius magnetic fields $B\\gg m^2 e^3$ (we are using Gauss system of units, where $e^2 = \\alpha = 1/137$; also in all formulas $\\hbar = c = 1$). It is natural to look for the atomic energy levels in such strong magnetic fields studing the influence of Coulomb potential on the electrons occupying Landau levels \\cite{Lan}. A strong magnetic field confines an electron in the transverse direction while in the longitudinal direction an electron is bound by the weak Coulomb field of a nucleus. The cigar-shape wave function of an electron is formed with transverse size which equals Larmour radius and longitudinal size which is by $\\ln(a_B^2/a_H^2) \\equiv \\ln(B/m^2 e^3)$ smaller than Bohr radius. The ground state energy is larger than Rydberg constant by the square of the same logarithm: $E_0 = -(me^4/2)\\ln^2(B/m^2 e^3)$. One can easily get this logarithmic factor from the fact that in one-dimensional Coulomb potential energy diverges logarithmically at small distances. The divergency is regularized  at the longitudinal distances which equals $a_H$, where one-dimensional motion converts to a three-dimensional one. Atomic levels in such strong magnetic fields were found numerically in \\cite{1} (see also \\cite{2, 22}).  Our purpose is to understand the behaviour of the energy levels with the growth of a magnetic field. The point is that at superstrong magnetic fields $B\\ga m^2/e^3$ the polarization operator insertions into the photon propagator induce the dynamic photon ``mass'' $m_\\gamma^2 \\approx e^3 B$ \\cite{3, 33}. One would expect that the photon mass should screen Coulomb potential and shift energies of the atomic levels found in tree approximation.  Dirac equation spectrum in a constant homogenious magnetic field looks like \\cite{4}: \\begin{equation} \\varepsilon_n^2 = m^2 + p_z^2 + (2n+1)eB + \\sigma eB \\;\\; , \\label{1} \\end{equation} where $n = 0,1,2, ...$, $\\sigma = \\pm 1$ and the field is directed along axis $z$.\\footnote{This spectrum with the substitution of $2n+1+\\sigma$ by $2j, \\,\\,j=0,1,2,... $ was found by I.I. Raby \\cite{IR}.} In the magnetic fields we are interested in $\\varepsilon_n \\ga m/e$, and electrons are ultrarelativistic. The only exception is the lowest Landau level (LLL) which has $n=0$, $\\sigma = -1$. The energy of LLL electron equals its mass for $p_z = 0$ and the consideration of the nonrelativistic electron motion along $z$ axis is selfconsistent. LLL is interesting both practically and theoretically. An analog of the critical electric field $E_{cr} = m^2/e$ is the magnetic field $B_0 = m^2/e = 4.4 \\cdot 10^{13}$ gauss. Two orders larger superstrong fields $B\\ga m^2/e^3$ can exist at special neutron stars named magnetars. The temperature of an outer magnetar layer is not enough to populate the excited Landau levels and one can observe the transitions among the states to which LLL is splitted at the electric field of the nucleus. Freezing of the ground state energy in the superstrong magnetic fields discussed in the paper leads to the upper bound on the spectra of photons radiated from magnetars.\\footnote{I am grateful to S.I. Blinnikov for the discussions of magnetar physics.} To study the stability of the huge magnetic fields \\cite{5} one should also know the energy of the ground state as a function of a field.  So we are studying the energies of the states to which LLL splits in the presence of an atomic nucleus.  Since the electron at LLL moves along $z$ axis we will study in section 2 QED at $D=2$: the behaviour of electrons in two-dimensional space-time. The coupling constant $g$ has dimension of mass, so Coulomb potential as a function of $|z|$ depends on two dimensionfull parameters: $g$  and electron mass $m$. We will obtain the approximate analytical formula for Coulomb potential in $d= D-1 = 1$ which takes into account the photon polarization operator. We will see that for large $g$ (or small $m$) $g \\gg m$ Coulomb potential is screened.  In section 3 we will consider the physical case, $D =4$ QED. The analog of the coupling constant squared $g^2$ in the real world is the product $e^3 B$. The polarization operator in the magnetic fields $B\\gg m^2/e$ at $k_\\parallel ^2(\\equiv k_z^2) \\ll eB$ practically coincides with the one obtained in section 2 \\cite{77}. Nevertheless the screening at large distances $|z| \\gg 1/m$ does not occur: at $|z| \\gg 1/m$ we get a purely Coulomb potential $\\Phi(z) = e/|z|$. The screening occurs at small distances, and its influence on the ground state energy is determined in section 4. The results similar to those presented in sections 3 and 4 were obtained in \\cite{55} with the help of the numerical calculations. ",
        "conclusions": "The photon polarization operator leads to modifications of the atomic energy levels. The famous example is its contribution to the Lamb shift, the difference of the energies  of $2s_{1/2}$ and $2p_{1/2}$ levels of hydrogen. They are numerically small loop corrections to the values of energies determined by the tree level potential. The role of the photon polarization diagram in the superstrong magnetic fields $B > m^2/e^3 = 6.2 \\cdot 10^{15}$ gauss is qualitatively different. It determines the behaviour of the ground state energy: the formula obtained at tree level becomes invalid and the growth of the coupling energy with $B$ terminates at $B\\approx m^2/e^3$. Screening of Coulomb potential should be more important for the energies of even excited states which are more sensitive to the shape of the potential at small distances \\cite{155}. Degeneracy of even and odd excited states in the limit $B\\Longrightarrow\\infty$ is not lifted by the screening. We study the analogy of the electric potential in $d=1$ QED with massive electrons and in $d=3$ QED in strong magnetic fields $B > B_0 = m^2/e$ which originates from the coincidence of the polarization operators in these cases. A simple analytical expression which equals the polarization operator with 10\\% accuracy enables us to obtain an approximate formula for the electric potential of the point charge in $d=1$ QED with massive fermions and asymptotics of the potential in $d=3$ QED. In $d=1$ QED for a coupling constant $g$ larger than a fermion mass $m$ a tree level formula is modified at $|z|> 1/g$. In $d=3$ QED a tree level formula is modified at the distances $1/m> |z| > 1/\\sqrt{e^3 B}$ while at large distances $|z|> 1/m$ Coulomb law is valid.  Analogous results for $D=4$ were obtained in \\cite{55}.  The other aspect of the Coulomb potential in the strong magnetic field is investigated  in paper \\cite{10}: it is supposed that fermions obtained their mass due to a magnetic field (dynamical fermion mass).  I am grateful to S.I. Blinnikov, V.A. Novikov, L.B. Okun, V.S. Popov, and A.V. Smilga for useful discussions and to A.I.Rez, who brought to my attention Phys. Rev. publication \\cite{55}.  This work was supported by the grants RFBR 08-02-00494, Nsh-4172.2010.2 and by the contract of the RF Ministry of Science and Education No. 02.740.11.5158."
    },
    "1005/1005.5376_arXiv.txt": {
        "abstract": "The Sun's supergranulation refers to a physical pattern covering the surface of the quiet Sun with a typical horizontal scale of approximately 30~000~kilometres and a lifetime of around 1.8~days. Its most noticeable observable signature is as a fluctuating velocity field of 360~m/s rms whose components are mostly horizontal. Supergranulation was discovered more than fifty years ago, however explaining why and how it originates still represents one of the main challenges of modern solar physics.  \\smallskip  A lot of work has been devoted to the subject over the years, but observational constraints, conceptual difficulties and numerical limitations have all concurred to prevent a detailed understanding of the supergranulation phenomenon so far. With the advent of 21st century supercomputing resources and the availability of unprecedented high-resolution observations of the Sun, a stage at which key progress can be made has now been reached. A unifying strategy between observations and modelling is more than ever required for this to be possible.  \\smallskip  The primary aim of this review is therefore to provide  readers with a detailed interdisciplinary description of past and current research on the problem, from the most elaborate observational strategies to recent theoretical and numerical modelling efforts that have all taken up the challenge of uncovering the origins of supergranulation. Throughout the text, we attempt to pick up the most robust findings so far, but we also outline the difficulties, limitations and open questions that the community has been confronted with over the years.  \\smallskip  In the light of the current understanding of the multiscale dynamics of the quiet photosphere, we finally suggest a tentative picture of supergranulation as a dynamical feature of turbulent magnetohydrodynamic convection in an extended spatial domain, with the aim of stimulating future research and discussions. ",
        "introduction": "The story of supergranulation really started in Oxford when Avril B. Hart reported in 1953 the existence of a ``noisy'' fluctuating velocity field on top of the mean rotation speed of the solar equator that she was measuring \\citep{hart54}. Actually, it is most probable that this ``noise\" was already detected as early as 1915 by \\cite{plaskett16}. Two years after her first detailed report, \\cite{hart56} confirmed the discovery and was able to give the first estimate\\epubtkFootnote{In solar physics, an appropriate length unit is the megametre (Mm), also 1000~km.} of 26~Mm for the typical horizontal length scale of these ``velocity fluctuations\" (sic). Supergranulation was further recognised as a characteristic feature of the whole surface of the quiet Sun (the regions of weak magnetic fields, which represent the most important part of the solar surface) after the seminal work of \\cite{leighton62}, who published the first Doppler images of the Sun (and also the first detection of the five minutes oscillations). This work was soon after completed by another important paper by \\cite{SL64} who showed, amongst other results, the intimate relation between supergranulation and the magnetic network of the quiet Sun. \\smallskip  It is remarkable that all the fundamentals of supergranulation have basically been uncovered over that 1954--1964 decade. Since then, progresses have been much less spectacular, especially on the theoretical side. This is certainly why supergranulation is still a fascinating subject. We are still wondering where it comes from, what its exact relation with magnetic fields is, if it is a universal feature of solar type stars, or of stellar surface convection, if it plays a role in the solar dynamo(s), etc. All these questions are pending fifty years after the discovery.  There are many reasons why the solar physics community has not yet managed to answer them, several of which are not actually specific to the supergranulation problem. Most of these reasons are described in detail in this review, but it is worth pointing out a few important issues here as an introduction. \\smallskip  For a long time only a limited set of observables and restricted time-records of the evolution of the supergranulation pattern have been available, may it be on short (24 consecutive hours) or long timescales (a solar cycle). This has somewhat hindered the study of the detailed spatial structure of supergranulation and the identification of the physical factors that affect it (buoyancy, stratification, magnetic fields, rotation). From the theoretical and numerical perspectives, on the other hand, the strongly nonlinear physical nature of magnetised thermal convection in the outer layers of the Sun makes it extremely difficult to come up with a simple, unique, verifiable physical model of the process (a similar problem arises in many subfields of solar physics, if not of astrophysics). Overall, these observational uncertainties and theoretical or numerical limitations have somewhat negatively interfered to prevent a rigorous solution to the supergranulation problem. The currently fairly obfuscated state of affairs may nevertheless get clarified in a near future, as the solar physics community is now armed with both high-resolution solar observatories such as Hinode\\epubtkFootnote{Hinode is led by the Japanese  Aerospace Exploration Agency (JAXA) in collaboration with NASA, the Science and Technology Facilities Council (STFC), and the European Space Agency (ESA). Hinode is a Japanese mission developed, launched and operated by ISAS/JAXA, in partnership with NAOJ, NASA and STFC (UK). Additional operational support is provided by ESA and NSC (Norway). The project website can be found at \\url{http://solarb.msfc.nasa.gov/}} and large supercomputers that allow for increasingly realistic numerical simulations of the complex solar surface flows. However, it is clear that stronger connections between theory, numerics and observations need to be established for the problem to be resolved.  \\smallskip  In this review, we would like to introduce readers to this subject by first describing the full range of past and current research activities pertaining to the solar supergranulation, from the breadth of historical and modern observational results to the most elaborate numerical models of supergranulation convection.  We particularly wish to provide a useful guide to the abundant literature related to that theme, to point out the important findings in the field, but also to stress the limitations and difficulties that have been encountered over the years in order to help overcome them. To this end, we attempt to discuss the already existing or possible connections between various pieces of research and try to identify some important questions whose answer may be crucial to understand how and why supergranulation originates.  \\smallskip  The review is divided into eight parts, including this introduction. The next two sections offers some introductory material on the physics of deep and surface convection in the Sun (Section~\\ref{generalconsiderations}) and a brief recap on small-scale flows,  namely granulation and mesogranulation (Section~\\ref{granule}). Section~\\ref{obs} is dedicated to a presentation of observational facts that have been collected on supergranulation. We then carry on with the discussion of existing theoretical models to explain the origin of supergranulation in Section~\\ref{theory}. In Section~\\ref{numerics}, we discuss related numerical experiments. Our current knowledge on supergranulation is summarised and commented in Section~\\ref{wrapup} for the convenience of hurried readers.  In the light of the present understanding of multiscale dynamics of the quiet photosphere, we finally suggest a tentative picture of supergranulation as a dynamical feature of turbulent magnetohydrodynamic (MHD) convection in an extended spatial domain, with the aim of stimulating future research and discussions. We notably propose several numerical and observational diagnostics that could help make important progress on the problem in the near future (Section~\\ref{discussion}).  \\smallskip  We tried to make the paper readable by all astrophysicists, assuming only little background in that field and trying to avoid as much as possible the solar physicists jargon, or to explain it when necessary. ",
        "conclusions": "} This long tour of the main observational, theoretical and numerical results on the problem of supergranulation being completed, we are now in a position to provide hurried (as well as less hurried) readers with a synthetic presentation of what has been learned so far on the supergranulation phenomenon and what are the current issues. The presentation of a more personal outlook and suggestions for future research is deferred to Section~\\ref{discussion}.  \\subsection{Observations} The Sun's supergranulation is a large-scale coherent pattern detected in the surface layers of the quiet Sun. The \\emph{impression} given by observations is that it is simply superimposed on a stochastic, highly nonlinear background smaller-scale flow pattern, the granulation. Characterising the supergranulation velocity pattern requires monitoring solar surface flows over long times, over wide fields of views, or over a large set of independent observations. The properties of the supergranulation velocity field can be summarised as follows:  \\begin{itemize} \\item the \\emph{length scale} of the supergranulation flow, as given by the kinetic energy power spectrum of the horizontal component of solar surface flows, is in the range of 20--70~Mm with a preferred scale of 36~Mm (Section~\\ref{length_sc}). These results come from both Dopplergrams \\citep{hathaway2000} and from granule tracking in wide-field high resolution image series \\citep{rieutord08}. The size of the field must be sufficiently large to secure the statistical convergence of the results.  \\item The typical \\emph{size} of supergranules, defined as coherent diverging flow cells at the solar surface, is in the range 10--30~Mm \\citep{hirzberger08}. The derived average size is sensitive to the method used to identify supergranules.  \\item The \\emph{kinetic energy excess} associated with \\emph{supergranulation} in the power spectrum of solar surface flows lies \\emph{on the large-scale side of the injection range of photospheric turbulence} located at the granulation scale (Section~\\ref{length_sc} and  Figure~\\ref{figure:spectra}).  \\item The most recent estimate of the \\emph{lifetime} of supergranules, based on the largest sample of supergranules collected so far, is 1.6$\\pm0.7$~day \\citep{hirzberger08}. The dispersion in the measurements of supergranules lifetimes is also fairly large (Section~\\ref{time_sc}).  \\item Rms \\emph{horizontal velocities} at supergranulation scale are of the order of 350~m/s, while rms \\emph{vertical velocities} are around 30~m/s \\citep{hathaway2002}. As velocities depend on the scales considered, the relation between amplitude and scale, namely the power spectrum, provides the most suitable observable to estimate the amplitude of the supergranulation velocity field (Section~\\ref{velo_sc}).  \\item Local helioseismology indicates that supergranules are \\emph{shallow structures} (Section~\\ref{obsdepth}), possibly not deeper than 5~Mm \\citep{sekii07}. The mean vertical profile of the supergranulation flow is not very well constrained at the moment. More precise determinations are definitely called for. \\end{itemize} Note that the foregoing determinations are not independent of each other, because velocity scales can be derived from the combination of length and time scales. Namely, 30~Mm divided by 1.7 day gives 205~m/s, which is in reasonable agreement with direct measurements of supergranulation-scale velocities. In our view, the computation of the power spectra of solar surface flows provides one of the most robust methods to make progress on the determination of these various quantities  in the future. Most notably, an accurate determination of the \\emph{vertical velocity spectrum} of vertical velocities in the supergranulation range is still lacking.  \\smallskip  Besides this set of typical scales associated with the supergranulation velocity pattern, several other observational signatures and properties of supergranulation have been studied:  \\begin{itemize} \\item \\emph{horizontal intensity fluctuations} at supergranulation scales are very faint (Section~\\ref{obsintensity}). The latest studies indicate that supergranules are slightly warmer at their centre. The temperature drop is less than 3~K \\citep{meunier07b,goldbaum_etal09}.  \\item Supergranulation is affected by the \\emph{global solar rotation} (Section~\\ref{obsrot}). Locally, supergranules are anticyclonic structures \\citep{gizon03a}, their mean vertical vorticity is negative in the northern hemisphere and positive in the southern one. The supergranulation pattern has been observed to propagate anisotropically in the prograde direction \\citep{gizon03b}.  \\item Supergranulation has \\emph{dynamical interactions} with the \\emph{magnetic fields} of the quiet Sun. Most notably, supergranules are strongly correlated with the magnetic network (Section~\\ref{obsmag}). Correlations between the size of supergranules and the strength of network and internetwork fields have been evidenced recently \\citep{meunier07a}. The solar-cycle dependence of the pattern remains uncertain though, as various papers have been giving contradicting results. \\end{itemize}  The bounds on intensity variations seem to be well established now. The proper rotation of supergranules, as measured by their local mean vertical vorticity, is also well constrained by local helioseismology. On the other hand, we believe that more work is required to quantitatively constrain the interaction of supergranules with magnetic fields. A determination of the \\emph{magnetic energy spectrum} of the quiet Sun over a wide range of scales would be extremely useful to put constraints on the physical processes at the origin of the network and internetwork fields and on their interactions with supergranulation (see Section~\\ref{discussion} below).  \\smallskip  Figure~\\ref{figure:super_pict} is an attempt to depict the standard view of the supergranulation phenomenon, as constrained by the observations summarised above.  \\epubtkImage{fig/super_pict.pdf}{ \\begin{figure}[ht] \\centerline{\\includegraphics[width=1.\\linewidth]{fig/super_pict}} \\caption[Schematic view of the supergranulation phenomenon, as constrained by observations.]{A schematic view of the supergranulation phenomenon, as constrained by observations. $\\lambda$ is the scale where the horizontal kinetic energy spectral density is maximum. $d$ is the diameter of ``coherent structures'' (supergranules). The red and blue patches depict the warm and cold regions of the flow. I.N.B denotes the internetwork magnetic field. Note that the indicated internetwork and network fields geometries roughly correspond to the standard historical picture of quiet Sun magnetic fields and their relation to supergranulation (Section~\\ref{obsmag}). As discussed in Sections~\\ref{obsINB} and \\ref{discussion}, this picture must be significantly nuanced in reality, as the dichotomy between network and internetwork fields is probably not quite as clear as indicated in this drawing.} \\label{figure:super_pict} \\end{figure} }  \\subsection{Theory} Two major types of physical scenarios have been suggested to explain the origin of supergranulation:  \\begin{itemize} \\item \\emph{thermal convection scenarios}, in which buoyancy is the main driver of the supergranulation flow (Section~\\ref{theoryconv}). Various effects (magnetic fields, shear, rotation, effective boundary conditions) have been explored within the framework of linear and weakly nonlinear theory to explain the size of of the supergranulation pattern, its weak thermal signature, its oscillations and propagation.  \\item \\emph{Collective interaction scenarios}, whereby supergranulation emerges as a large-scale coherent pattern triggered by nonlinear interactions of vigorous smaller-scale structures like granulation (see \\cite{rieutord2000} and Section~\\ref{theoryLSinst}). These scenarios have mostly been explored quantitatively through ``toy model'' simulations \\citep{rast03b,crouch07} that do not incorporate the full complexity of dynamical MHD equations. In our view, direct numerical simulations provide the  most promising way of making progress on this side in the future. \\end{itemize} A shared property of all models is the looseness of the approximations on which they rely (e.g. linear theory with turbulent viscosity parametrisation, or purely phenomenological arguments on the nature of dynamical interactions between granules and their potential large-scale instabilities). Completely distinct theoretical arguments can easily be tuned to produce results that are all broadly consistent with observations. This degeneracy makes it impossible to discriminate between various scenarios and to come up with a proper theoretical explanation for the origin of supergranulation that could be unambiguously validated by observations.  \\smallskip  Finally, it is possible but certainly not obvious that supergranulation can be explained quantitatively by a simple mathematical theoretical model. In any case, one of the most urgent tasks to overcome some of the previously mentioned shortcomings is to figure out if the basic assumptions and arguments on which current theoretical models rely  (linear theory, effective boundary conditions, convection in uniform magnetic fields etc.) are justified, and to test them quantitatively with the help of large-scale numerical simulations.  \\subsection{Numerical simulations} The complexity and nonlinearity of the physical environment of supergranulation is extraordinary: vigorous turbulent small-scale flows in a strongly stratified atmosphere, ionisation physics, rotation, shear and tortuous magnetic fields geometries at all observable scales may all have something to do with the supergranulation phenomenon. As argued several times in this review, numerical simulations have a unique potential for approaching this complexity. They have now become an unavoidable tool to uncover the real nature of supergranulation and to test the various qualitative theoretical pictures described in Section~\\ref{theory}.  \\smallskip  Numerical simulations dedicated to the supergranulation problem are still in their infancy though, mostly because they remain awfully expensive in terms of computing time. The latest generation of numerical experiments, summarised in Table~\\ref{table:performances}, barely accommodates for the scale of supergranulation. The main results obtained so far are summarised below:  \\begin{itemize} \\item \\emph{global spherical simulations} \\citep{derosa02} exhibit a supergranulation-like pattern, but the scale of this pattern is dangerously close to the grid scale of the simulations (Section~\\ref{simglobal}).  \\item In \\emph{local large-scale idealised simulations} \\citep{cattaneo01,rincon05}, two patterns can be singled out of the continuum of turbulent scales: a granulation pattern forming in the upper thermal boundary layer, and a larger-scale, extremely energetic mesoscale pattern, which extends through the whole convective layer (Figure~\\ref{figure:ideal}). Whether or not this pattern has anything to do with supergranulation or with  the hypothetical solar mesogranulation is not understood (see Section~\\ref{simlargelocal} for an in-depth discussion).  \\item \\emph{Local large-scale realistic simulations of hydrodynamic convection} \\citep{stein09} do not exhibit any significant energy excess at supergranulation scales in spite of the presence of Hydrogen and Helium ionizations in the model (Section~\\ref{simlargelocal2}). This result therefore tends to disprove the ``classical'' \\cite{SL64} supergranulation theory.   \\item \\emph{Local large-scale realistic simulations of MHD convection} reveal the formation of a magnetic network at scales ranging from mesoscales to supergranulation scales (Section~\\ref{simmhd}). What sets the scale of this network and the emergence of supergranulation as a special scale in these simulations has not been investigated yet, but a recent study \\citep{ustyugov09} suggests that strong magnetic flux concentrations play a significant role in the scale-selection process. \\end{itemize}  \\smallskip  Numericists will have to address several important issues in the forthcoming years. One of the main problems is that all dedicated simulations to date are still fairly dissipative (much more than the Rayleigh-B\\'enard simulations described in Section~\\ref{numconv}, for instance). Local large-scale simulations, for instance, barely accommodate 10 grid points within a granule. This kind of resolution is not sufficient to capture all the dynamics of solar surface flows, as the viscous and magnetic dissipation scales are both much smaller than 100~km (Section~\\ref{turbulentscales}) at the solar surface. As mentioned in Section~\\ref{realvsideal} and \\ref{numconc}, resolving dissipation scales properly has recently turned out to be essential to make progress on several turbulent MHD problems, such as magnetic field generation (dynamo action) by non-helical turbulent velocity fields. A related point is that uncovering the full dynamical physics of large scales and avoiding spurious finite-box effects requires both very large numerical domains and large integration times of the simulations, which is not ensured in today's experiments. This point is easily illustrated by the supergranulation-scale dichotomy between global and local simulations discussed in Section~\\ref{simlarge}.  \\smallskip  Overall, the current computing limitations are such that numerical simulations are still far away from the parameter regime typical of the Sun. Hence, one cannot exclude that all simulations to date miss some critical multiscale dynamical phenomena, either purely hydrodynamic or MHD. Large-scale simulations are also currently too expensive for any decent scan of the parameter space of the problem to be possible. However, it is fair to say that the perspective of petaflop computations holds the promise of significant numerical breakthroughs in a ten-years future."
    },
    "1005/1005.2283_arXiv.txt": {
        "abstract": "{The Sun is a magnetic star whose cyclic activity is thought to be linked to internal dynamo mechanisms. A combination of numerical modelling with various levels of complexity is an efficient and accurate tool to investigate such intricate dynamical processes. } {We investigate the role of the magnetic buoyancy process in 2D Babcock-Leighton dynamo models, by modelling more accurately the surface source term for poloidal field.}{ To do so, we reintroduce in mean-field models the results of full 3D MHD calculations of the non-linear evolution of a rising flux tube in a convective shell. More specifically, the Babcock-Leighton source term is modified to take into account the delay introduced by the rise time of the toroidal structures from the base of the convection zone to the solar surface.} {We find that the time delays introduced in the equations produce large temporal modulation of the cycle amplitude even when strong and thus rapidly rising flux tubes are considered. Aperiodic modulations of the solar cycle appear after a sequence of period doubling bifurcations typical of non-linear systems. The strong effects introduced even by small delays is found to be due to the dependence of the delays on the magnetic field strength at the base of the convection zone, the modulation being much less when time delays remain constant. We do not find any significant influence on the cycle period except when the delays are made artificially strong.}{ A possible new origin of the solar cycle variability is here revealed. This modulated activity and the resulting butterfly diagram are then more compatible with observations than what the standard Babcock-Leighton model produces.} ",
        "introduction": "Our Sun is a prime example of a very turbulent and magnetically active star. Its robust 22-yr activity cycle originates in the periodic polarity reversal of the Sun's internal large-scale magnetic field. As a result, sunspots emerge at the solar surface with statistically well-defined dynamical and morphological characteristics. These sunspots are thought to be the surface manifestations of strong toroidal magnetic structures created at the base of the convection zone and which rise through the plasma under the effect of magnetic buoyancy (\\cite{Parker55}). This toroidal field undergoes cyclic variations along with the poloidal field, which flips polarity at sunspot cycle maxima. Even if a robust regular activity is easily exhibited in the Sun, a significant modulation of both the amplitude and the frequency of the cycle has been observed. In particular, periods of strongly reduced activity have been revealed, the most famous being the Maunder minimum between 1650 and 1700, during when no sunspots were to be seen at the solar surface (\\cite{Eddy76}). Today, understanding the origins of such a variability has become crucial. Not only is the impact on satellites, astronauts, power grids or radio communications significant but the climatologists community now tends to address more and more the question of the influence of solar variability on the Earth climate (e.g. \\cite{Bard06}, \\cite{Lean05}).  The classical explanation for the cyclic activity of the large-scale magnetic field is that a dynamo process acts in the solar interior to regenerate the three components of the magnetic field and sustain them against ohmic dissipation. The inductive action of the complex fluid motions would thus be responsible for the vigorous regeneration of magnetic fields and for its non-linear evolution in the solar interior (see \\cite{Charbonneau051} and \\cite{Miesch05} for recent reviews on the subject).    Understanding how these complex physical processes operating in the solar turbulent plasma non-linearly interact is very challenging. One successful and powerful approach is to rely on multi-dimensional magnetohydrodynamic (MHD) simulations. In this context, two types of numerical experiments have been performed since the 70's: kinematic mean-field axisymmetric dynamo models which solve only the mean induction equation (\\cite{Steenbeck69, Roberts72, Stix76, Krause80}) and full 3D global models which explicitly solve the full set of MHD equations (\\cite{Gilman83, Glatzmaier85, Cattaneo99, Brun04}). Those two approaches are complementary and needed since 2D mean-field models are limited by the fact that they rely on simplified descriptions of complex physical processes such as turbulence and since the cost of 3D models make it difficult, as of today, to provide any reliable predictions concerning the large-scale magnetic cycle.  As far as the first kind of numerical simulations is concerned, a particular model has been favoured by a part of the community, namely the Babcock-Leighton flux-transport model first proposed by \\cite{Babcock61} and \\cite{Leighton69}. This model has proved to be very efficient at reproducing several properties of the solar cycle as the equatorward branch of activity or the phase relationship between toroidal and poloidal fields. It has thus been extensively used since the 90's (\\cite{Wang91, Durney95, Dikpati99, Nandy01}), even to make tentative predictions of the next solar cycle (\\cite{Dikpati06}). In this formulation, the poloidal field owes its origin to the tilt of active regions emerging at the Sun's surface at various latitudes during the solar cycle. It thus relies on a different mechanism than ``classical'' $\\alpha\\Omega$ dynamo models for which the poloidal field is generated by turbulent helical motions twisting and shearing toroidal field lines inside the convection zone. The emergence of tilted active regions at the photosphere is the result of the complex non-linear evolution of strong toroidal structures from the base of the convection zone (CZ), where they are created through the $\\Omega$-effect. It is thus natural to rely on 3D MHD simulations of rising toroidal structures in the convection zone to gain some insight on the best way to model the Babcock-Leighton source term.   Many models carried out since the 80's relied on the assumption that toroidal flux is organised in the form of discrete flux tubes which will rise cohesively from the base of the CZ up to the solar surface. These models enabled to demonstrate that the initial strength of magnetic field was an important parameter in the evolution of the tube and that the active regions tilts could be explained by the action of the Coriolis force on the magnetic structure (\\cite{DSilva93, Fan94, Caligari95}). More sophisticated multidimensional models were then developed and extended to the upper part of the CZ and the transition to the solar atmosphere (\\cite{Magara04, Archontis05, Cheung07,  Martinez08}). Computations were then performed to study the influence of convective turbulent flows on the dynamical evolution of flux ropes inside the CZ (\\cite{Dorch01, Cline03, Fan03}, Jouve $\\&$ Brun 2009). In particular in \\cite{Jouve09}, such computations were made in a global spherical geometry and thus allowed to assess the combined role of hoop stresses, Coriolis force, convective plumes, turbulence, mean flows and sphericity on the tube evolution and on the subsequent emerging regions, along with the usual parameters such as field strength, twist of the field lines or magnetic diffusion. How can the results of 3D simulations of this presumed major step of the dynamo loop be reintroduced in 2D mean-field models? This is the question we are addressing in this work. A particular point we retain from these simulations is that the rise velocity and thus the rise time of magnetic structures is strongly dependent on the initial field strength at the base of the convection zone. Indeed, magnetic buoyancy competes with convective flows to control the trajectory and the speed of the flux tubes while they rise. Since the Babcock-Leighton source term relies on this process to regenerate poloidal fields, it may be worth trying to introduce magnetic-field dependent time delays caused by rising toroidal fields in those models and analyze the influence on the solar cycle.  Indeed, time delays built up in solar dynamo models have been shown to cause long-term modulation of the dynamo cycle and under some circumstances to lead to a chaotic behaviour (\\cite{Yoshimura78}, \\cite{Charbonneau052} or \\cite{Wilmot05}). In the framework of mean-field dynamo theory, several possibilities have been studied to explain the variability of the solar cycle. They mainly fell into two categories: stochastic forcing or dynamical nonlinearities. Indeed, as we mentioned above, the solar convection zone is highly turbulent and it would be surprising if the dynamo processes acting inside this turbulent plasma were nicely regular. The influence of stochastic fluctuations in the mean-field dynamo coefficients has been studied in various models (e.g. \\cite{Hoyng88,Ossendrijver96,Weiss00,Charbonneau00}). Moreover, the dynamical feedback of the strong dynamo-generated magnetic fields is likely to be significant enough to produce non-linear effects on the activity cycle. A number of models have introduced these non-linear effects (\\cite{Proctor77, Tobias97, Moss00,Bushby06,Rempel06}) and have resulted in the production of grand minima-like periods or other strong modulation of the cyclic activity. However, time delays have hardly been considered and if they were, they were mainly due to the advection time by meridional flow. Indeed, the time-scale of the buoyant rise of flux tubes was considered to be so small compared to the cycle period (and to the meridional flow turnover time) that this particular step was assumed to be instantaneous. However, we would like to address the question of the influence of magnetic field dependent delays on the cycle produced by Babcock-Leighton models and especially on its potential modulation. We propose to do so in the present paper.  This article is organized as follows: Section \\ref{sect_3D} summarizes the results of 3D calculations which will be used in our modified 2D mean-field Babcock-Leighton model. The formulation of this new model is then presented, with a particular focus on the modification of the surface source term. Results of 2D models are shown in Sect. \\ref{sect_2D} and analyzed in the following two sections. Sect. \\ref{sect_6th} and \\ref{sect_5th} present and study the behaviour of a reduced set of ordinary differential equations designed to gain some insight on the results of the 2D model. We discuss the results and conclude in Sect. \\ref{sect_conclu}.           The idea of this work is to design a new 2D  mean-field flux transport Babcock-Leighton model which takes into account the findings of 3D calculations concerning the rise of strong toroidal structures from the base of the convection zone up to the surface. In this section, we thus present the usual mean-field equations and the basic ingredients used in the model, with a particular focus on the new Babcock-Leighton surface source term.   \\subsection{The mean field Babcock-Leighton model}  To model the solar global dynamo, we use the hydromagnetic induction equation, governing the evolution of the magnetic field ${\\bf B}$ in response to advection by a flow field ${\\bf V}$ and resistive dissipation.  \\begin{equation} \\frac{\\partial {\\bf B}}{\\partial t}=\\nabla\\times ({\\bf V} \\times{\\bf B})-\\nabla\\times(\\eta\\nabla\\times{\\bf B}) \\end{equation} where $\\eta$ is the magnetic diffusivity.  As we are working in the framework of mean-field theory, we express both magnetic and velocity fields as a sum of large-scale (that will correspond to mean field) and small-scale (associated with fluid turbulence) contributions. Averaging over some suitably chosen intermediate scale makes it possible to write two distinct induction equations for the mean and the fluctuating parts of the magnetic field. A closure relation is then used to express the electromotive force in terms of mean magnetic field, leading to a simplified mean-field equation. In this work we will replace the emf by a surface Babcock-Leighton term (Babcock 1961; Leighton 1969; Wang et al. 1991; Dikpati \\& Charbonneau 1999; Jouve \\& Brun 2007a) as described in details below.         We work here in spherical geometry and under the assumption of axisymmetry. Reintroducing a poloidal/toroidal decomposition of the magnetic and velocity fields in the mean induction equation, we get two coupled partial differential equations, one involving the poloidal potential $A_{\\phi}$ and the other concerning the toroidal field $B_{\\phi}$.  \\begin{eqnarray} \\label{eqA2} \\frac{\\partial {A_{\\phi}}}{\\partial t}&=&\\frac{\\eta}{\\eta_{t}} (\\nabla^{2}-\\frac{1}{\\varpi^{2}})A_{\\phi}- R_{e}\\frac{\\bf{v}_{p}}{\\varpi}\\cdot\\nabla(\\varpi A_{\\phi}) \\nonumber \\\\ &+&C_{s}S(r,\\theta,B_{\\phi}) \\end{eqnarray}  \\begin{eqnarray} \\label{eqB2} \\frac{\\partial {B_{\\phi}}}{\\partial t}&=&\\frac{\\eta}{\\eta_{t}} (\\nabla^{2}-\\frac{1}{\\varpi^{2}})B_{\\phi} +\\frac{1}{\\varpi}\\frac{\\partial(\\varpi B_{\\phi})}{\\partial r}\\frac{\\partial (\\eta/\\eta_{t})}{\\partial r} \\nonumber  \\\\ &-&R_{e}\\varpi {\\bf v}_{p}\\cdot\\nabla(\\frac{B_{\\phi}}{\\varpi})-R_{e}B_{\\phi}\\nabla\\cdot{\\bf v}_{p} \\nonumber  \\\\ &+&C_{\\Omega}\\varpi(\\nabla\\times(\\varpi A_{\\phi}{\\bf \\hat{e}}_{\\phi}))\\cdot\\nabla\\Omega \\end{eqnarray}  \\noindent where $\\varpi=r\\sin\\theta$, $\\eta_{t}$ is the turbulent magnetic diffusivity (diffusivity in the convective zone), ${\\bf v}_{p}$ the flow in the meridional plane (i.e. the meridional circulation), $\\Omega$ the differential rotation, $S(r,\\theta,B_{\\phi})$ the Babcock-Leighton source term for poloidal field. In order to write these equations in a dimensionless form, we choose as length scale the solar radius ($R_{\\odot}$) and as time scale the diffusion time ($R_{\\odot}^2/\\eta_{t}$) based on the envelope diffusivity ($\\eta_{t}$). This leads to the appearance of three control parameters $C_{\\Omega}=\\Omega_{0}R_{\\odot}^2/\\eta_{t}$, $C_{s}=s_{0}R_{\\odot}/\\eta_{t}$ and $R_{e}=v_{0}R_{\\odot}/\\eta_{t}$ where $\\Omega_{0}, s_{0}, v_{0}$ are respectively the rotation rate and the typical amplitude of the surface source term and of the meridional flow.  Equations $\\ref{eqA2}$ and $\\ref{eqB2}$ are solved in an annular meridional cut with the colatitude $\\theta$ $\\in [0,\\pi]$ and the radius (in dimensionless units) $r \\in [0.6,1]$ i.e from slightly below the tachocline ($r=0.7$) up to the solar surface, using the STELEM code. This code has been thoroughly tested and validated thanks to an international mean field dynamo benchmark involving 8 different codes (\\cite{Jouve08}). At $\\theta=0$ and $\\theta=\\pi$ boundaries, both $A_{\\phi}$ and $B_{\\phi}$ are set to 0. Both $A_{\\phi}$ and $B_{\\phi}$ are set to $0$ at $r=0.6$. At the upper boundary,  we smoothly match our solution to an external potential field, i.e. we have vacuum for $r \\geq 1$. As initial conditions we are setting a confined dipolar field configuration, i.e the poloidal field is set to $\\sin\\theta / r^{2}$ in the convective zone and to $0$ below the tachocline whereas the toroidal field is set to $0$ everywhere. \\\\  \\subsection{The standard physical ingredients and the modified surface source term}  We now need to prescribe the amplitude and profile of the various ingredients acting in equations $\\ref{eqA2}$ and $\\ref{eqB2}$, namely the rotation, the diffusivity, the meridional flow and the Babcock-Leighton source term.  The rotation profile captures some realistic aspects of the Sun's angular velocity, deduced from helioseismic inversions (Thompson et al. 2003), assuming a solid rotation below $0.66$ and a differential rotation above the interface.  \\begin{eqnarray} {\\Omega(r,\\theta)}&=&\\Omega_{c}+\\frac{1}{2}[1+\\rm erf(2\\frac{r-r_{c}}{d_{1}})] \\nonumber \\\\ & &(\\Omega_{Eq}+a_{2}\\cos^2\\theta+a_{4}\\cos^4\\theta-\\Omega_{c}) \\end{eqnarray}  \\noindent with $\\Omega_{Eq}=1$,  $\\Omega_{c}= 0.93944$,  $r_{c}=0.7$, $d_{1}=0.05$, $a_{2}=-0.136076$ and $a_{4}=-0.145713$ and where $\\rm erf$ is the error function. With this profile, the radial shear is maximal at the tachocline.    We assume that the diffusivity in the envelope $\\eta$ is dominated by its turbulent contribution whereas in the stable interior $\\eta_{\\rm c} \\ll \\eta_{\\rm t}$. We smoothly match the two different constant values with an error function which enables us to quickly and continuously transit from $\\eta_{\\rm c}$ to $\\eta_{\\rm t}$ i.e.   \\begin{equation} \\eta(r)=\\eta_{\\rm c}+\\frac{1}{2}\\left(\\eta_{\\rm t}-\\eta_{\\rm c}\\right)\\left[1+{\\rm erf}\\left(\\frac{r-r_{\\rm c}}{d}\\right)\\right], \\label{eqeta} \\end{equation}  \\noindent with  ${\\eta_{\\rm c}}=10^9 \\,\\rm cm^2\\rm s^{-1}$ and $d=0.03$.  In Babcock-Leighton flux-transport dynamo models, meridional circulation is used to link the two sources of the magnetic field namely the base of the CZ and the solar surface. For all the models presented in this paper we use a large single cell per hemisphere, directed poleward at the surface, vanishing at the bottom boundary $r=0.6$ and thus penetrating a little below the tachocline. We take a stream function  \\begin{equation} \\psi(r,\\theta)=-\\frac{2}{\\pi}\\frac{(r-r_{b})^2}{(1-r_{b})}\\sin\\left(\\pi\\frac{r-r_{b}}{1-r_{b}}\\right)\\cos\\theta\\sin\\theta, \\label{eqpsi} \\end{equation}  \\noindent which gives, through the relation  ${\\bf u_{\\rm p}}=\\nabla \\times(\\psi \\hat {\\bf e}_{\\phi})$, the following components of the meridional flow  \\begin{eqnarray} u_{r}&=&-\\frac{2(1-r_{\\rm b})}{\\pi r}\\frac{(r-r_{\\rm b})^2}{(1-r_{\\rm b})^2} \\sin\\left(\\pi\\frac{r-r_{\\rm b}}{1-r_{\\rm b}}\\right)(3\\cos^2\\theta-1), \\end{eqnarray} \\begin{eqnarray} u_{\\theta}&=&\\Bigg[\\frac{3r-r_{\\rm b}}{1-r_{\\rm b}} \\sin\\left(\\pi\\frac{r-r_{\\rm b}}{1-r_{\\rm b}}\\right)+\\frac{r\\pi}{1-r_{\\rm b}}\\frac{(r-r_{\\rm b})}{(1-r_{\\rm b})} \\cos\\left(\\pi\\frac{r-r_{\\rm b}}{1-r_{\\rm b}}\\right)\\Bigg] \\nonumber \\\\ & \\times &\\frac{2(1-r_{\\rm b})}{\\pi r}\\frac{(r-r_{\\rm b})}{(1-r_{\\rm b})}\\cos\\theta\\sin\\theta, \\end{eqnarray} \\noindent with $r_{\\rm b}=0.6$.    In Babcock-Leighton dynamo models, the poloidal field owes its origin to the tilt of magnetic loops emerging at the solar surface. Since these emerging loops are thought to rise from the base of the convection zone through magnetic buoyancy, we see that we can directly relate the way we model the Babcock-Leighton (BL) source term and the results of 3D calculations shown in Sect \\ref{sect_3D}.     In the standard model, the source term is confined in a thin layer at the surface and is made to be antisymmetric with respect to the equator, due to the sign of the Coriolis force which changes from one hemisphere to the other. These features are retained in our modified version.  However, the standard term is proportional to the toroidal field at the base of the convection zone at the same time, implying an instantaneous rise of the flux tubes from the base to the surface where they create tilted active regions. The 3D calculations showed that the rise velocity and thus the rise time depend on the field strength at the base of the CZ, we thus introduce in our modified version of the source term, a magnetic field-dependent time delay in the toroidal field at the base of the convection zone. We thus take into account the time delay between the formation of strong toroidal structures at the base of the convection zone and the surface regeneration of poloidal field.   The modified expression of the source term is thus  \\begin{eqnarray} S(r,\\theta,B_{\\phi})&=& \\frac{1}{2}\\Bigg[1+\\rm erf(\\frac{r-r_{2}}{d_{2}})\\Bigg]\\Bigg[ 1-erf(\\frac{r-1}{d_{2}})\\Bigg] \\nonumber \\\\ & \\times & \\Bigg[1+({\\frac{B_{\\phi}(r_{c},\\theta,t-\\tau_B)}{B_{0}}})^{2}\\Bigg]^{-1} \\nonumber \\\\ & \\times &\\cos\\theta \\sin\\theta \\,\\, B_{\\phi}(r_{c},\\theta,t-\\tau_B) \\label{eq_s} \\end{eqnarray}  where $r_{2}=0.95$, $d_{2}=0.01$, $B_{0}=10^4 \\, \\rm G$, with the time delay $\\tau_B$ proportional to the inverse of the magnetic energy at the base of the convection zone, namely $$ \\tau_B(\\theta,t) = \\tau_0/B_{\\phi}(r_c,\\theta,t)^2 $$  In our 3D simulations, the approximate rise time for a $6 \\times 10^5 \\, \\rm G$ field is indeed about 4 times that of a $3 \\times 10^5 \\, \\rm G$ field. This takes into account the more significant effects of convective downdrafts and Ohmic diffusion in the ``weak B'' case. Note that the time delay is then dependent both on space and time.   Finally, we take into account the fact that very strong toroidal structures (more than $10^5 \\, \\rm G$) are not influenced by the Coriolis force enough to gain a significant tilt as they reach the surface. To do so, we introduce a quenching term in the surface source which will make the regeneration of the poloidal field less effective when the toroidal field at the base of the CZ is strong enough. Inversely, when the flux tubes are too weak, they will not be able to reach the surface at all, and will not take part into the regeneration term for the poloidal field. % We thus prevent the weakest flux tubes (or equivalently the most delayed ones) to take part in the regeneration of poloidal fields at the surface. To do so, the source term is set to zero when the delays are above a certain threshold value corresponding to a rise time of approximately half a solar cycle. ",
        "conclusions": "\\label{sect_conclu}  In this work, we have introduced a more physically accurate model for the poloidal field source term in Babcock-Leighton flux transport dynamo models. To do so, some results from 3D MHD calculations of rising flux tubes from the base of the convection zone to the surface were reintroduced in 2D mean-field models. In particular, the rise time of flux ropes, which was assumed to be negligible in previous simulations, was here taken into account by introducing magnetic energy-dependent time delays in the BL source term. Through the use of a combination of analytical and numerical techniques applied to an adapted reduced system of non linear equations, we show that the system exhibits a sequence of bifurcations leading to a chaotic behaviour. The time evolution of the solution at that stage is qualitatively very similar to the full 2D model and more importantly to the actual solar activity, with strong modulation especially on the cycle amplitude.    Assuming the rise of flux tubes from the base of the convection zone and the surface to be instantaneous and thus the time delays due to magnetic buoyancy to be negligible can sound reasonable since these delays are weak compared to the other time-scales of the system. In particular, we can expect the time delay due to the advection by the meridional flow (which has a characteristic time-scale of a few solar cycles, see for example \\cite{Charbonneau00}) to be much more effective on the behaviour of the dynamo-generated magnetic field. As a consequence, it is mainly the effect of meridional flow which has been tested in previous models (e.g. \\cite{Wilmot05}). However, we showed here that the time delays due to flux tubes rising have in fact a significant effect to modulate the cycle on time scales very large compared to the delays. The main reason for this behaviour is that the delays are not fixed but depend on the modulated toroidal energy itself. To illustrate this explanation, we show on Figure \\ref{figure_bq_mod} the time evolution of toroidal fields (delayed and undelayed) for a case similar to what we showed in the core of the paper (a $\\vert B \\vert^2$-dependent delay with $\\tau_0=3.4 \\times 10^{-2}$, i.e. just after the Hopf bifurcation) and the same evolution when the delay is fixed to the mean value reached by the delay in the previous run.      \\begin{figure}[h!] \\centering \\includegraphics[width=8.5cm]{./fig_bq_modul.ps} \\includegraphics[width=8.5cm]{./fig_bq_cst.ps} \\caption{Time evolution of the toroidal fields amplitude (undelayed in plain line and delayed in dotted line) for a delay dependent on $\\vert B \\vert^2$ just after the Hopf bifurcation (upper panel) and for a delay fixed to the mean value of the previous run (lower panel).} \\label{figure_bq_mod} \\end{figure}   On this figure, both the basic toroidal field and the delayed field are shown, stressing the time shift existing between the two. The striking result is that a long-term modulation appears only in the case of a varying delay. In the fixed-delay case, although the time-shift between the basic toroidal field and its delayed counterpart is quite significant compared to the cycle period, no modulation is created. On the contrary, in the varying-delay case, even though the strongest fields are almost not affected by the time-shift, a strong modulation on the amplitude is built up. As a result, the fact that small delays can affect the long-term evolution of the dynamo cycle seems to be linked to the variability of these delays and more specifically to their dependence on the magnetic field strength. Since magnetic buoyancy is much more efficient in the solar interior for strong toroidal field structures, these kind of delays are likely to appear in reality. Consequently, a new origin of the solar cycle modulation may have been revealed here, thanks to the input of 3D MHD models into 2D mean-field dynamo calculations.  Only a part of the results of 3D calculations were reintroduced here but other effects can also be taken into account. In particular, hoop stresses and the Coriolis force are known since the first thin flux tubes simulations to deflect the trajectory of buoyantly rising magnetic fields inside the convection zone, this effect being stronger when the fields are less intense. This feature could also be introduced in BL models by modifying the source term for poloidal field. In particular, the source term is likely to get a contribution from the same latitude a the initial position of the flux tubes when those are strong enough and from a lower latitude when the flux tubes have been sufficiently influenced by the Coriolis force and the hoop stresses to rise parallel to the rotation axis. This may result in a different shape for the butterfly diagram with slightly more activity at higher latitudes, provided that weak magnetic structures are assumed to be able to reach the surface and create active regions. This could lead to reconsider the correspondence which is often made between toroidal magnetic field generated in 2D dynamo models at the base of the convection zone and the actual sunspot migration observed at the solar surface.  The variations of the solar cycle period are also known to be significant. For instance, cycle 23 has been significantly longer than the previous ones, with a duration of about 13 years (1996-2009). In the models computed in this work, the modulation of the solar activity is particularly visible on the cycle amplitude and is much less obvious on the period. However, the influence of rising flux tubes on the meridional flow was not taken into account here. Since the cycle period of the solutions resulting from these flux-transport models is very sensitive to the amplitude of the meridional circulation, a modification of the flow by rising flux tubes is very likely to modify the frequency of the dynamo solution and thus to also introduce a modulation on the cycle period. This feature is currently being investigated.  A considerable step forward would obviously be to develop a self-consistent global model with buoyant toroidal structures built up and making their way from the base of the convection zone where they become unstable to the photosphere where they create active regions. Unfortunately, this has not been achieved yet due to numerous physical and numerical difficulties. In the mean time, we have shown here that inputs from 3D MHD models simulating a particular step of the dynamo cycle can significantly improve 2D mean-field calculations. In particular, we have shown that they can even help reproducing a feature of the solar cycle (its variability) which was not present in the standard model before. We will continue to develop these ideas in future work."
    },
    "1005/1005.2626_arXiv.txt": {
        "abstract": "s{ The Extragalactic Background Light (EBL) from the infrared (IR) through the ultraviolet (UV) is dominated by emission from stars, either directly or through absorption and reradiation by dust.  It can thus give information on the star formation history of the universe.  However, it is difficult to measure directly due to foreground radiation fields from the Galaxy and solar system.  Gamma-rays from extragalactic sources at cosmological distances (blazars and gamma-ray bursts) interact with EBL photons creating electron-positron pairs, absorbing the gamma-rays.  Given the intrinsic gamma-ray spectrum of a source and its redshift, the EBL can in principle be measured.  However, the intrinsic gamma-ray spectra of blazars and GRBs can vary considerably from source to source and the from the same source over short timescales.  A maximum intrinsic spectrum can be assumed from theoretical grounds, to give upper limits on the EBL absorption from blazars at low redshift with very high energy (VHE) gamma-ray observations with ground-based Atmospheric Cherenkov telescopes.  The Fermi-LAT observations of blazars and GRBs can probe EBL absorption at higher redshifts.  The lower energy portion of the LAT spectrum of these sources is unattenuated by the EBL, so that extrapolating this to higher energies can give the maximum intrinsic spectrum for a source.  Comparing this to the observed higher energy LAT spectrum will then give upper limits on the EBL absorption.  For blazars which have been detected by both the Fermi-LAT and at higher energies by Cherenkov telescopes, combined LAT-VHE observations can put more stringent constraints on the low redshift EBL.  These procedures above can also be reversed: for sources with an unknown redshift, a given EBL model and gamma-ray spectrum can lead to an upper limit on the source's redshift.  } ",
        "introduction": "\\label{intro}  The night sky appears dark to the naked eye, but in fact glows faintly in the IR through the optical and UV.  At these wavelengths, the background light is dominated by emission from the atmosphere, solar system, and Milky Way.  There is also a much smaller extragalactic component from all of the stars which have ever existed, through direct emission (in the UV-optical) and through absorption and reradiation by dust (in the IR).  Due to the weakness of this extragalactic background light (EBL) to other components, direct measurement of the EBL is extremely difficult~\\cite{bernstein02,mattila03,bernstein07}.  The other background components can be avoided to some extent by using instruments on spacecraft which have left the atmosphere~\\cite{bernstein02,hauser98} or the solar system~\\cite{toller83,edelstein00}.  However, it is unlikely that spacecraft will leave our Galaxy in the near future, so uncertainties in direct measurements will remain.  Number counts in the IR and optical can be used to find EBL lower limits~\\cite{madau00,marsden09}, as discussed by Beelen and Penin in these proceedings. Modeling~\\cite{kneiske04,stecker06,franceschini08,gilmore09,finke10} has been an important tool for constraining the EBL intensity and tying it to basic astrophysics such as the star formation rate density, dust absorption, initial mass function, cosmological expansion rate, and others.  Fig.\\ \\ref{EBLmodels_fig} shows many EBL measurements, constraints and models, and Hauser \\& Dwek~\\cite{hauser01} present a thorough review.    \\begin{figure} \\begin{center} \\includegraphics[scale=0.34]{EBLmodels_06_stretch} \\caption{EBL models, measurements, and constraints.  See Finke {\\em et al.} for details and references. \\label{EBLmodels_fig}} \\end{center} \\end{figure}   The EBL photons interact with $\\g$-rays from cosmological sources to produce $e^+e^-$ pairs, absorbing the $\\g$-rays so that the observed flux $F_{obs}(E) = F_{int}(E)\\exp[-\\tau_{\\g\\g}(E)]$ where $F_{int}(E)$ is the unabsorbed source flux as a function of observed energy $E$, and $\\tau_{\\g\\g}(E)$ is the EBL absorption optical depth.  If $F_{int}(E)$ is known, a measurement of the observed $\\g$-ray spectrum from these sources can be used to probe the EBL.  The intrinsic spectrum is not generally known, however it is possible to determine an upper limit either from theory or from extrapolating a lower energy, unattenuated spectrum to higher energies.  This is discussed further in the next sections.  From the upper limit on $F_{int}(E)$ and the measurement of $F_{obs}(E)$ with a $\\g$-ray telescope, an upper limit on $\\tau_{\\g\\g}(E)$ can be calculated and compared to theoretical predictions. ",
        "conclusions": ""
    },
    "1005/1005.4620_arXiv.txt": {
        "abstract": "{This paper discusses some of the challenges of spectro-polarimetric observations with a large aperture solar telescope such as the ATST or the EST. The observer needs to reach a compromise among spatial and spectral resolution, time cadence, and signal-to-noise ratio, as only three of those four parameters can be pushed to the limit. Tunable filters and grating spectrographs provide a natural compromise as the former are more suitable for high-spatial resolution observations while the latter are a better choice when one needs to work with many wavelengths at full spectral resolution. Given the requirements for the new science targeted by these facilities, it is important that 1)tunable filters have some multi-wavelength capability; and 2)grating spectrographs have some 2D field of view. } ",
        "introduction": "The Advanced Technology Solar Telescope (ATST), the European Solar Telescope (EST) and the so-called plan B for the space mission Solar-C are all being designed with multi-wavelength spectro-polarimetry as one of the top priorities. This kind of observations will permit novel studies of the processes taking place through the solar atmosphere, from the photosphere to the corona, as they traverse very different physical regimes with conditions that typically vary over several orders of magnitude.  Observing the chromosphere/corona at disparate wavelengths with spectro-polarimetry is still challenging even with the large-aperture telescopes currently under development. The science goals dictate requirements that typically include the following: \\begin{itemize} \\item High spatial resolution. We are interested in observing processes that take place at very small spatial scales. Thus, the telescope needs to be able to observe near the diffraction limit. \\item High cadence. The demand for high spatial resolution imposes also requirements on the time cadence. If a particular feature is moving with a projected velocity $v$, integrating over time scales comparable or longer than $t={\\Delta x}/v$ (where $\\Delta x$ is the spatial resolution) will result in image degradation. The sound speed (which is often taken as a first estimate for $v$) in the photosphere is typically around 6~km~s$^{-1}$. However, in the chromosphere it is significantly larger and one often finds supersonic flows. \\item High spectral resolution. Many science goals require not only imaging but also the observation of detailed line profiles to be able to track shock waves in the shape of chromospheric lines, as well as other interesting phenomena that produce emission features or other imprints on the intensity or polarization spectral profiles. Ideally one would want to resolve the Doppler width of the lines. Many interesting chromospheric lines are from H and He ions and have large Doppler widths, but others (particularly Ca) impose requirements for a better resolution. \\item High signal-to-noise ratio. Polarimetric signals are often weak, especially those coming from the chromosphere. Typically, a continuum signal-to-noise ratio of 1000 is required for diagnostics based on the Zeeman effect (some times an order of magnitude better is required for Hanle-effect studies). \\end{itemize}  Unfortunately, it is not possible to meet all four of these requirements simultaneously, regardless of the telescope size. Observers will have to make compromises and optimize the instrument configuration for the problem at hand. Compromising will then be a key notion in observational solar physics for the next decade. ",
        "conclusions": "The tremendous leap in spatial resolution that will be provided by the ATST and the EST will require careful rethinking of most conventional observing modes, particularly long slit spectroscopy. Without at least some limited 2D capability, it is not clear how to make these observations work in practice, particularly with the requirement of seamless multi-wavelength operation that is so vital to the science goals of these facilities. It should be noted, though, that a large aperture is still useful for moderate- or even low-resolution observations in which one needs to reach very high polarimetric sensitivities (e.g., for some Hanle effect observations).  Efficient exploitation of the new large solar telescopes with minimum downtimes will be of paramount importance, especially since it now appears that they will concentrate most of the resources devoted by the solar community to observations. This condition implies shifting the telescope operations to a new model more similar to what is currently employed in the large night-time facilities or the space-borne instrumentation such as Hinode. However, there is the risk that the community will be left without a facility where new experimental ideas can be tested. It is extremely important that the new operational model for the ATST/EST considers a provision of at least a small fraction of the observing time to test new ideas or generally speaking, high risk/high return potential campaigns."
    },
    "1005/1005.4599_arXiv.txt": {
        "abstract": "We consider the interpretation of the MiniBooNE low-energy anomaly and the Gallium radioactive source experiments anomaly in terms of short-baseline electron neutrino disappearance in the framework of 3+1 four-neutrino mixing schemes. The separate fits of MiniBooNE and Gallium data are highly compatible, with close best-fit values of the effective oscillation parameters $\\Delta{m}^2$ and $\\sin^2 2\\vartheta$. The combined fit gives $\\Delta{m}^2 \\gtrsim 0.1 \\, \\text{eV}^2$ and $ 0.11 \\leq \\sin^2 2\\vartheta \\leq 0.48 $ at $2\\sigma$. We consider also the data of the Bugey and Chooz reactor antineutrino oscillation experiments and the limits on the effective electron antineutrino mass in $\\beta$-decay obtained in the Mainz and Troitsk Tritium experiments. The fit of the data of these experiments limits the value of $\\sin^2 2\\vartheta$ below 0.10 at $2\\sigma$. Considering the tension between the neutrino MiniBooNE and Gallium data and the antineutrino reactor and Tritium data as a statistical fluctuation, we perform a combined fit which gives $\\Delta{m}^2 \\simeq 2 \\, \\text{eV}$ and $ 0.01 \\leq \\sin^2 2\\vartheta \\leq 0.13 $ at $2\\sigma$. Assuming a hierarchy of masses $ m_{1}, m_{2}, m_{3} \\ll m_{4} $, the predicted contributions of $m_{4}$ to the effective neutrino masses in $\\beta$-decay and neutrinoless double-$\\beta$-decay are, respectively, between about 0.06 and 0.49 and between about 0.003 and 0.07 eV at $2\\sigma$. We also consider the possibility of reconciling the tension between the neutrino MiniBooNE and Gallium data and the antineutrino reactor and Tritium data with different mixings in the neutrino and antineutrino sectors. We find a $2.6\\sigma$ indication of a mixing angle asymmetry. ",
        "introduction": "Introduction}  Neutrino oscillations have been observed in solar, atmospheric and long-baseline reactor and accelerator experiments. The data of these experiments are well fitted in the framework of three-neutrino mixing, in which the three flavor neutrinos $\\nu_{e}$, $\\nu_{\\mu}$, $\\nu_{\\tau}$ are unitary linear combinations of three massive neutrinos $\\nu_{1}$, $\\nu_{2}$, $\\nu_{3}$ with the solar (SOL) and atmospheric (ATM) squared-mass differences \\begin{align} \\null & \\null \\Delta{m}^2_{21} = \\Delta{m}^2_{\\text{SOL}} \\simeq 8 \\times 10^{-5} \\, \\text{eV}^2 \\,, \\label{002} \\\\ \\null & \\null |\\Delta{m}^2_{31}| \\simeq |\\Delta{m}^2_{32}| = \\Delta{m}^2_{\\text{ATM}} \\simeq 2 \\times 10^{-3} \\, \\text{eV}^2 \\,, \\label{003} \\end{align} where $\\Delta{m}^2_{jk} = m_{j}^2 - m_{k}^2$ and $m_{j}$ is the mass of the neutrino $\\nu_{j}$ (see Refs.~\\cite{hep-ph/9812360,hep-ph/0211462,hep-ph/0310238,hep-ph/0405172,hep-ph/0506083,hep-ph/0606054,GonzalezGarcia:2007ib,Giunti-Kim-2007}).  Besides these well-established observations of neutrino oscillations, there are at least three anomalies which could be signals of short-baseline neutrino oscillations generated by a larger squared-mass difference: the LSND $\\bar\\nu_{\\mu}\\to\\bar\\nu_{e}$ signal \\cite{hep-ex/0104049}, the Gallium radioactive source experiments anomaly \\cite{0901.2200,1001.2731}, and the MiniBooNE low-energy anomaly \\cite{0812.2243}. In this paper we consider the MiniBooNE and Gallium anomalies, which can be explained by short-baseline electron neutrino disappearance \\cite{0707.4593,0711.4222,0902.1992} in the effective framework of four-neutrino mixing, as explained in Sections~\\ref{009} and \\ref{019}. On the other hand, the LSND anomaly is disfavored by the results of the MiniBooNE $\\nu_{\\mu}\\to\\nu_{e}$ experiment \\cite{0704.1500,0812.2243} and may require another explanation \\cite{hep-ph/0010308,0705.0107,0706.1462,0710.2985,0805.2098,0906.1997,0906.5072}.  In Refs.~\\cite{0707.4593,0902.1992} we proposed to explain the MiniBooNE low-energy anomaly \\cite{0704.1500,0812.2243} through the disappearance of electron neutrinos due to very-short-baseline oscillations into sterile neutrinos generated by a squared-mass difference $ \\Delta{m}^2 $ larger than about $20\\,\\text{eV}^2$. In that case, the analysis of the MiniBooNE data is simplified by the fact that the effective survival probability $P_{\\nu_{e}\\to\\nu_{e}}$ is practically constant in the MiniBooNE energy range from 200 to 3000 MeV. In this paper we extend the analysis of MiniBooNE data to lower values of $ \\Delta{m}^2 $, considering the resulting energy dependence of the effective short-baseline (SBL) electron neutrino and antineutrino survival probability \\begin{equation} P_{\\boss{\\nu}{e}\\to\\boss{\\nu}{e}}^{\\text{SBL}}(L,E) = 1 - \\sin^2 2\\vartheta \\sin^2\\!\\left( \\frac{ \\Delta{m}^2 L }{ 4 E } \\right) \\,, \\label{004} \\end{equation} where $L$ is the neutrino path length and $E$ is the neutrino energy (CPT invariance implies that the survival probabilities of neutrinos and antineutrinos are equal; see Ref.~\\cite{Giunti-Kim-2007}).  The two-neutrino-like effective short-baseline survival probability in Eq.~(\\ref{004}) is obtained in four-neutrino schemes (see Refs.~\\cite{hep-ph/9812360,hep-ph/0405172,hep-ph/0606054,GonzalezGarcia:2007ib}), which are the simplest extension of three-neutrino mixing schemes which can accommodate the two small solar and atmospheric squared-mass differences in Eqs.~(\\ref{002}) and (\\ref{003}), and one larger squared-mass difference for short-baseline neutrino oscillations, \\begin{equation} |\\Delta{m}^2_{41}| = \\Delta{m}^2 \\gtrsim 0.1 \\, \\text{eV}^2 \\,. \\label{005} \\end{equation} The existence of a fourth massive neutrino corresponds, in the flavor basis, to the existence of a sterile neutrino $\\nu_{s}$.  In this paper we consider 3+1 four-neutrino schemes, since 2+2 four-neutrino schemes are disfavored by the combined constraints on active-sterile transitions in solar and atmospheric neutrino experiments \\cite{hep-ph/0405172}. For simplicity, we consider only 3+1 four-neutrino schemes with \\begin{equation} m_{1}, m_{2}, m_{3} \\ll m_{4} \\,, \\label{006} \\end{equation} which give the $\\Delta{m}^2_{41}$ in Eq.~(\\ref{005}) and appear to be more natural than the other possible 3+1 four-neutrino schemes in which either three neutrinos or all four neutrinos are almost degenerate at a mass scale larger than $\\sqrt{\\Delta{m}^2}$ (see Refs.~\\cite{hep-ph/9812360,hep-ph/0405172,hep-ph/0606054,GonzalezGarcia:2007ib}).  In 3+1 four-neutrino schemes the effective mixing angle in the effective short-baseline electron neutrino survival probability in Eq.~(\\ref{004}) is given by (see Refs.~\\cite{hep-ph/9812360,hep-ph/0405172,hep-ph/0606054,GonzalezGarcia:2007ib}) \\begin{equation} \\sin^2 2\\vartheta = 4 |U_{e4}|^2 \\left( 1 - |U_{e4}|^2 \\right) \\,. \\label{007} \\end{equation} In this paper we assume that the value of $|U_{\\mu4}|^2$ is so small that the effective short-baseline muon neutrino survival probability is practically equal to unity and short-baseline $\\boss{\\nu}{\\mu}\\leftrightarrows\\boss{\\nu}{e}$ are negligible\\footnote{ In 3+1 four-neutrino schemes the effective short-baseline muon neutrino survival probability has the form in Eq.~(\\ref{004}) with $\\sin^2 2\\vartheta$ replaced by $\\sin^2 2\\vartheta_{\\mu\\mu} = 4 |U_{\\mu4}|^2 \\left( 1 - |U_{\\mu4}|^2 \\right)$. The effective short-baseline $\\boss{\\nu}{\\mu}\\leftrightarrows\\boss{\\nu}{e}$ transition probability is given by $ P_{\\boss{\\nu}{\\mu}\\leftrightarrows\\boss{\\nu}{e}}^{\\text{SBL}}(L,E) = \\sin^2 2\\vartheta_{e\\mu} \\sin^2\\!\\left( \\frac{ \\Delta{m}^2 L }{ 4 E } \\right) $, with $\\sin^2 2\\vartheta_{e\\mu} = 4 |U_{e4}|^2 |U_{\\mu4}|^2$ (see Refs.~\\cite{hep-ph/9812360,hep-ph/0405172,hep-ph/0606054,GonzalezGarcia:2007ib}). }. This assumption is justified by the lack of any indication of $\\nu_{\\mu}\\to\\nu_{e}$ transitions in the MiniBooNE experiment \\cite{0704.1500,0812.2243} and the limits on short-baseline muon neutrino disappearance found in the CDHSW \\cite{Dydak:1984zq}, CCFR \\cite{Stockdale:1985ce} and MiniBooNE \\cite{0903.2465} experiments. We do not consider the MiniBooNE antineutrino data \\cite{0904.1958}, which have at present statistical uncertainties which are too large to constraint new physics \\cite{0902.1992}.  The plan of the paper is as follows. In Section~\\ref{009} we discuss the analysis of MiniBooNE data. In Section~\\ref{019} we present an update of the analysis of Gallium data published in Ref.~\\cite{0711.4222} and the combined analysis of MiniBooNE and Gallium data. In Section~\\ref{027} we discuss the implications of the measurements of the effective electron neutrino mass in Tritium $\\beta$-decay experiments and their combination with reactor neutrino oscillation data. In Section~\\ref{041} we present the results of the combined analysis of MiniBooNE, Gallium, reactor and Tritium data and in Section~\\ref{045} we present the corresponding predictions for the effective masses measured in $\\beta$-decay and neutrinoless double-$\\beta$-decay experiments. In Section~\\ref{058} we calculate the mixing angle asymmetry between the neutrino and antineutrino sectors which could explain the tension between the neutrino and antineutrino data under our short-baseline $\\nu_{e}$-disappearance hypothesis. In Section~\\ref{060} we draw the conclusions.  \\begin{figure*} \\begin{center} \\begin{tabular}{lr} \\includegraphics*[bb=33 147 578 695, width=0.48\\linewidth]{fig-01a.eps} & \\includegraphics*[bb=33 147 578 695, width=0.48\\linewidth]{fig-01b.eps} \\end{tabular} \\end{center} \\caption{ \\label{008} Expected number of $\\nu_{e}$ events compared with MiniBooNE data, represented by the black points. The energy bins are numbered with the index $j$. The uncertainty is represented by the vertical error bars, which represent the sum of statistical and uncorrelated systematic uncertainties. (a) Expected number of $\\nu_{e}$-like events $N_{\\nu_{e},j}^{\\text{cal}}$ calculated by the MiniBooNE collaboration. $N_{\\nu_{e},j}^{\\text{cal}}$ is given by the sum of the $\\nu_{e}$-induced events ($N_{\\nu_{e},j}^{\\nu_{e},\\text{cal}}$) and the misidentified $\\nu_{\\mu}$-induced events ($N_{\\nu_{e},j}^{\\nu_{\\mu},\\text{cal}}$). (b) Best-fit value of the number of $\\nu_{e}$-like events $N_{\\nu_{e},j}^{\\text{the}}$ obtained with the hypothesis of $\\nu_{e}$ disappearance. $N_{\\nu_{e},j}^{\\text{the}}$ is given by the sum of $ N_{\\nu_{e},j}^{\\nu_{e},\\text{the}} = f_{\\nu} P_{\\nu_{e}\\to\\nu_{e}}^{(j)} N_{\\nu_{e},j}^{\\nu_{e},\\text{cal}} $ and $ N_{\\nu_{e},j}^{\\nu_{\\mu},\\text{the}} = f_{\\nu} N_{\\nu_{\\mu}}^{\\text{cal}} $. The best-fit values of $f_{\\nu}$, $\\sin^2 2\\vartheta$ and $\\Delta{m}^2$ are those in the first column of Tab.~\\ref{017} (MB$\\nu$). } \\end{figure*} ",
        "conclusions": ""
    },
    "1005/1005.2189_arXiv.txt": {
        "abstract": "{We have carried out two extremely deep surveys with SPIRE, one of the two cameras on {\\it Herschel}, at 250\\,$\\mu$m, close to the peak of the far-infrared background. We have used the results to investigate the evolution of the rest-frame 250-$\\mu$m luminosity function out to $z =\\rm 2$. We find evidence for strong evolution out to $z \\simeq\\rm 1$ but evidence for at most weak evolution beyond this redshift. Our results suggest that a significant part of the stars and metals in the universe today were formed at $z \\preceq\\rm 1.4$ in spiral galaxies.} ",
        "introduction": "The discovery that approximately half the energy ever radiated by galaxies is received on Earth in the far-infrared waveband (Puget et al.\\ 1996; Fixsen et al.\\ 1998) implies that galaxies must show strong evolution that is hidden from optical telescopes (Gispert, Lagache \\& Puget 2000). A decade ago, deep surveys with the SCUBA camera on the James Clerk Maxwell Telescope resolved much of the far-infrared background (FIRB) at 850 $\\mu$m into individual sources (Barger et al.\\ 1998; Hughes et al.\\ 1998). These sources are mostly extremely luminous dust-enshrouded galaxies at $\\rm z > 2$ (Chapman et al.\\ 2005) with an average implied star-formation rate (if the ultimate source of the energy is star formation) of $\\simeq$400 $\\rm M_{\\odot}\\ year^{-1}$ (Coppin et al.\\ 2006), much greater than the star-formation rates in galaxies like our own.  However, the energy density in the FIRB at 850 $\\mu$m is $\\simeq$30 times less than at 200 $\\mu$m where the FIRB is at a maximum, and both the spectral shape of the FIRB and statistical `stacking' analyses (Dole et al. 2006; Pascale et al.\\ 2009) imply that much of the FIRB is actually produced by sources at lower redshift (Gispert et al.\\ 2000; Dole et al. 2006; Pascale et al.\\ 2009). The launch of the Herschel Space Observatory (Pilbratt et al.\\ 2010) in May 2009 has given us the opportunity to resolve a significant fraction of the FIRB at wavelengths where its energy density is at a maximum. In this letter, using the first data from the Herschel Multi-tiered Extragalactic Survey (HerMES; Oliver et al., in prep), we investigate the evolution implied by the existence of the FIRB by measuring the evolution of the galaxy luminosity function at 250 $\\mu$m. We everywhere assume a standard concordance cosmology: $\\rm \\Omega_M=0.28,\\ \\Omega_{\\Lambda}=0.72,\\ H_0 = 72\\ km\\ s^{-1}\\ Mpc^{-1}$. ",
        "conclusions": "On the assumption that the star-formation rate is proportional to the rest-frame 250-$\\mu$m luminosity, the lack of evidence for any strong evolution at $\\rm z > 1$ is consistent with investigations that have concluded that the overall star-formation rate per unit comoving volume was approximately constant at $\\rm z > 1$ (Steidel et al.\\ 1999; Gispert et al.\\ 2000; Hopkins 2004). We estimate from the GOODS-North catalogue that we have resolved $\\simeq$20\\% of the FIRB at 250 $\\mu$m. Fig.~3 shows that the sources making up this top 20\\% of the FIRB are at moderate redshift ($z \\sim 1$). Although we cannot say anything directly about the sources responsible for the missing 80\\% of the FIRB, the stacking analysis of Pascale et al.\\ (2009) from BLAST data suggests that the remainder of the FIRB at this wavelength is also from sources at moderate redshift.  A revealing way to look at these early HerMES results is to use the relationship between the production of metals and the background radiation associated with this metal production (Peacock 1993). Using a value for the integrated FIRB of $\\rm 14\\ nW\\ sr^{-1}$ (Fixsen et al.\\ 1998), we obtain a relationship between the mass of metals and the fraction, $\\epsilon$, of the FIRB produced in a particular redshift interval: \\smallskip $$ M_{metals} = 9 \\times 10^6 \\epsilon (1+z)\\ M_{\\odot}\\ Mpc^{-3}. $$ \\smallskip \\noindent Fig.~5 shows the metals produced as a function of redshift on the assumption that the redshift distributions of GOODS-North and LH-North are representative of the FIRB as a whole. The figure suggests that most of the metals (and therefore most of the massive stars) formed at a moderate redshift ($z \\preceq 1.4$). This conclusion is consistent with the overall star-formation rate in the universe being constant at $\\rm z > 1$, and the apparent decline of metal production at high redshift is simply the consequence of the relationship between cosmic time and redshift. However, the decline in the figure at high redshifts is relatively small, and since we are making the large assumptions that the redshift distributions of our samples are characteristic of the FIRB as a whole, it is possible that observations that resolve more of the FIRB and at several wavelengths rather than a single wavelength will modify this conclusion.  The excellent images that exist from the Hubble Space Telescope for the GOODS-North {\\it Herschel} sources mean that we can examine the nature of the galaxies responsible for this metal production. We have used the z-band images, since at $\\rm z \\simeq 1$ this band corresponds in the rest-frame to the B-band. Of the 26 galaxies in the redshift interval $\\rm 0.8 < z < 1.2$, 12 show clear signs of spiral structure, and there may well be more if some of the galaxies have spiral arms that are below the surface-brightness detection threshold of the images. Thus the galaxies found in the first HerMES images appear to be quite different from the major mergers found in the SCUBA surveys (Ivison et al.\\ 2000; Tacconi et al.\\ 2008). Studies of galaxy evolution based on {\\it Spitzer} 24-$\\mu$m surveys have also found that the {\\it Spitzer} galaxies at $\\rm z \\simeq 1$ are often spirals rather than major mergers (Elbaz et al.\\ 2007) but the crucial advance that {\\it Herschel} has made possible is that we can determine the cosmological importance of these objects, in terms of the total masses of stars and metals formed in them.  \\begin{figure} \\centering \\includegraphics[width=6cm]{14675fg5.eps} \\caption{An estimate of the total metals formed per cubic Mpc in different redshift intervals using the method described in the text. We have made the estimates from the redshift distributions of GOODS-North (red) and LH-North (green) on the assumption that these redshift distributions represent the FIRB as a whole.} \\end{figure}"
    },
    "1005/1005.2971_arXiv.txt": {
        "abstract": "In addition to the Sun, six other stars are known to harbor multiple planets and debris disks: HD 69830, HD 38529, HD 128311, HD 202206, HD 82943 and HR 8799. In this paper we set constraints on the location of the dust-producing planetesimals around the latter four systems. We use a radiative transfer model to analyze the spectral energy distributions of the dust disks (including two new $Spitzer$ $IRS$ spectra presented in this paper), and a dynamical model to assess the long-term stability of the planetesimals'  orbits.  As members of a small group of stars that show evidence of harboring a  multiple planets and planetesimals, their study can help us learn  about the diversity of planetary systems. ",
        "introduction": "Surveys with the {\\it Spitzer} Space Telescope have been spectacularly successful at identifying infrared excess emission associated with planetary debris disks around A- through K-type stars. The excesses at 70 $\\mu$m are associated with cool dust located at distances from the stars analogous to the position of the Kuiper Belt (KB) in the solar system (Moro-Mart\\'in et al. 2008), although much larger numbers of objects must lie in these exo-Kuiper belts to account for the detected level of emission (Trilling et al. 2008; Hillenbrand et al. 2008; Carpenter et al. 2009). The most sensitive studies find 70 $\\mu$m excesses around $\\sim$ 15-20\\% of mature solar-type stars over the entire 10 Myr to 10 Gyr age range (Trilling et al. 2008). The excesses at 24$\\mu$m are generally associated with warmer dust and disappear relatively quickly as the host star ages; about 30\\% of the stars of the Pleiades age ($\\sim$ 120 Myr) show excess emission (Sierchio et al. 2010) whereas by the age of Praesepe ($\\sim$ 600 Myr), the 24 $\\mu$m excesses have almost completely disappeared (Meyer et al. 2008; G\\'asp\\'ar et al. 2009). Because the expected lifetimes of the debris dust grains are much shorter than the ages of the stars, it is inferred that the dust originates from collisional activity in reservoirs of planetesimals left over from the planet formation process (hence the term {\\it debris} dust). To sustain the dust production, it is necessary that large planetesimals (1000 km-sized) or unseen planets stir the planetesimals so they continue to collide with each other. Planets are also responsible for constraining the planetesimals in some zones and clearing them from others, thus determining much of the structure of the debris system. The way the patterns of debris disks activity decay with age is consistent with the expectation that the inner zones of a planetary system have relatively short dynamical time scales, whereas dynamical activity unfolds slowly at the distance of the Kuiper Belt.  A highlight of the recent surveys is the first detection of debris disks around stars with planets (Beichman et al. 2005; Moro-Mart\\'in et al. 2007a). These planets orbit within several AU of their parent star, whereas the cold dust emitting the observed far-IR radiation generally resides tens of AU away. Despite the separation between the dust and planets, it is still possible  for the planet to shape the structure of  the dust (and planetesimals) disk. In the multiple-planet system HD 38529, for example, secular resonances excited by planets at 0.13 and 3.7 AU create regions at tens of AUs that are unstable for orbiting planetesimals (Moro-Mart\\'in et al. 2007b).  Several additional systems have been identified as having both multiple planets (capable of exciting secular resonances) and orbiting debris (indicating the presence of planetesimals; Bryden et al. 2009, Su et al. 2009). Presumably there are many more stars with debris-disk excesses that also harbor multiple but undiscovered planets. Detailed studies of the known examples can reveal aspects of their behavior that help us understand the diversity of planetary systems. In this paper we study four of these systems: HD 128311, HD 202206, HD 82943 and HR 8799. In Section \\ref{obs} we describe the planet and debris dust detections for each one of these systems, and present new $Spitzer$ $IRS$  detections  for the debris disks around HD 202206 and HD 82943. In Section \\ref{sed} we use a radiative transfer model to identify the range of parameters (dust mass and dust location) that would fit the observed spectral energy distribution (SED). Due to the high fractional luminosity, the grain-grain collisional time-scale is shorter than the Poynting-Robertson (P-R) time-scale for all these systems and therefore we expect the dust to trace the location of the dust-producing planetesimals.  In Section \\ref{dyn}, we use a dynamical model to assess the long-term orbital stability of the putative dust-producing planetesimals, taking into account the effect of secular resonances.  Putting together the results from the SED and dynamical analysis, in Section \\ref{loc} we discuss the potential location of the dust-producing planetesimal belts. ",
        "conclusions": "In this paper we have studied the possible planet-planetesimal configurations of four multi-planet systems, of which three are radial-velocity systems -- HD 128311, HD 202206, HD 82943 -- and one is a directly imaged system -- HR 8799. We have quantified where the zone of influence lies of planets on the dust-producing planetesimals: for HR 8799 it extends to nearly 20 AU from the orbit of the outermost planet, while for the three radial-velocity systems it extends to about 4 AU.  A previous paper that studied HD38529, another radial-velocity multi-planet system, found that the influence of the planets in this case extends out to $\\sim$ 10 AU (determining the inner edge of the disk), and becomes dominant again at a $\\sim$ 55 AU (due to a secular resonance that probably determines the outer edge of the dust disk (Moro-Mart\\'in et al. 2007b). We conclude that radial-velocity multi-planet systems generally have zones of influence within a few to ten AU; more precise determinations will require individual modeling of a system. The influence can be extended much further through secular resonances and similar behavior.  For the three radial-velocity multi-planet systems studied in this paper, we have constructed fits to the spectral energy distributions of the debris disks. If we use astronomical silicates with a size distribution down to the blow-out size, the emitting regions of the disks must be so far from the star that they are well outside the zones of influence of the radial velocity planets. This behavior depends critically on the optical properties of the grains (size distribution and optical constants of the grain material). If we were to adopt optical constants typical of an ice-silicate mixture instead of astronomical silicates, the dust may be located closer to the star. In all three cases, we find that single-size 10 $\\mu$m astronomical silicates reproduce well the observed SED, and could lie at the edge of the zone of influence of the planets and therefore the disk would be sculpted by them.   \\begin{center} {\\it Acknowledgments} \\end{center} We thank Hal Levison for providing  skeel-SyMBA  for the dynamical simulations, Sebastian Wolf for providing DDS for the SED models, Alexander Krivov for his careful reading of the manuscript and useful comments and Daniel Fabrycky and Ruth Murray-Clay for useful comments. This work is based on observations made with the {\\it Spitzer} Space Telescope, which is operated by the Jet Propulsion Laboratory, managed for NASA by the  California Institute of Technology. A.M.M. acknowledges funding from the Spanish MICINN (Ram\\'on y Cajal Program and grants AYA2009-07304 and CONSOLIDER INGENIO 2010CSD2009-00038), the Michelson Fellowship and the {\\it Spitzer} archival grant 40412. She thanks the Isaac Newton Institute for Mathematical Sciences at Cambridge University for support. RM acknowledges support from grants by NSF (AST-0806828) and NASA (NNX08AQ65G).  \\newpage"
    },
    "1005/1005.0373_arXiv.txt": {
        "abstract": "{We investigate the possibility of using the ratio between the 2-10 keV flux and the [Ne~V]3426 emission line flux (X/NeV) as a diagnostic diagram to discover heavily obscured, possibly Compton-Thick Active Galactic Nuclei (AGN) in the distant Universe. While being on average about one order of magnitude fainter than the more commonly used [O~III]5007 emission line, the [Ne~V]3426 line can be observed with optical spectroscopy up to $z\\sim 1.5$, whereas the [O III]5007 line is redshifted out of the optical bands already at $z\\sim0.8$. First, we calibrate a relation between X/NeV and the cold absorbing column density $N_H$ using a sample of 74 bright, nearby Seyferts with both X-ray and [Ne V] data available in the literature, and for which the column density is determined unambiguously. Similarly to what is found for the X-ray to [O III]5007 flux ratio (X/OIII), we found that the X/NeV ratio decreases towards large column densities, as expected if [Ne V]3426 emission is a good tracer of the AGN intrinsic power. Essentially all local Seyferts with X/NeV values below 15 are found to be Compton-Thick objects. At X/NeV values below 100, the percentage of Compton-Thick nuclei decreases to $\\sim 50\\%$, but still $\\sim 80\\%$ of the considered sample is absorbed with $N_H>10^{23}$ cm$^{-2}$. Second, we apply this diagnostic diagram to different samples of distant obscured and unobscured QSOs in the Sloan Digital Sky Survey (SDSS). SDSS blue, unobscured, type-1 QSOs in the redshift range $z=[0.1-1.5]$ indeed show X/NeV values typical of unobscured Seyfert 1s in the local Universe. Conversely, SDSS type-2 QSOs at $z\\sim 0.5$ classified either as Compton-Thick or Compton-Thin on the basis of their X/OIII ratio, would have been mostly classified in the same way based on the X/NeV ratio. We apply the X/NeV diagnostic diagram to 9 SDSS obscured QSOs in the redshift range $z=[0.85-1.31]$, selected by means of their prominent [Ne~V]3426 line (rest $EW>4$\\AA) and observed with \\chandra\\ ACIS-S for 10ks each (8 of them as part of our proprietary program). Based on the X/NeV ratio, complemented by X-ray spectral analysis, 2 objects appear good Compton-Thick QSO candidates, 4 objects appear as Compton-Thin QSOs, while 3 have an ambiguous classification. When excluding from the sample broad lined QSOs with a red continuum and thus considering only genuine narrow-line objects, the efficiency in selecting Compton-Thick QSOs through the [Ne~V] line is about 50\\% (with large errors, though), more similar to what is achieved with [O~III] selection. We discuss the possibility of applying the X/NeV diagnostic to deep X-ray surveys to search for Compton-Thick Seyferts at $z\\sim 1$, i.e. those objects which are thought to be responsible for the ``missing\" X-ray background.  Finally, we compare the optical spectral properties of [Ne~V]-selected QSOs with those of other SDSS populations of obscured and unobscured QSOs. By restricting the analysis to objects in the same redshift (and luminosity) range z=[0.4-1.5], we found evidence that, at any given [Ne~V] luminosity, increasing obscuration is accompanied by increasing [O~II]3727 emission. This correlation is interpreted as evidence for enhanced star formation in obscured QSOs, which is consistent with current popular scenarios of BH-galaxy coevolution. ",
        "introduction": "\\label{introduction}  While the cosmological evolution of unobscured QSOs has been traced up to $z\\sim 6$, the evolution of obscured AGN is much more uncertain and is the subject of intense debate. The {\\it observed} number statistics in current AGN samples is dominated by unobscured objects which are easier to discover (e.g. $>10^5$ QSOs have been identified in the Sloan Digital Sky Survey), but several arguments suggest that obscured AGN must be intrinsically more numerous. Deep X-ray surveys (see \\citealt{bh05} for a review) have indeed shown that, towards faint X-ray fluxes, the surface density of obscured AGN overtakes that of unobscured AGN. Also, population synthesis models of the cosmic X-ray background (XRB), suggest that obscured AGN outnumber unobscured ones by a factor which ranges from $\\sim 2$ to $\\sim 8$, depending on the considered luminosity regime (\\citealt{gch07}; see \\citealt{tuv09} and \\citealt{balla06} for a steeper luminosity dependence). To understand the cosmological history of accretion onto supermassive black holes (SMBHs) it is therefore necessary to map and understand the population of obscured AGN and, in particular, of the most obscured and hence elusive ones, the so-called Compton-Thick (CT) nuclei, i.e. those obscured by column densities above $\\sim 10^{24}$ cm$^{-2}$. The population of moderately obscured AGN in fact does not completely account for the XRB peak intensity at 30 keV, to which CT AGN are expected to contribute from $\\sim 10\\%$ \\citep{tuv09} to $\\sim 25-30\\%$ \\citep{gch07}, depending on the XRB model assumptions. In addition, the presence of a large population of CT AGN across the cosmic epochs, would help in reconciling the measured mass function of local SMBHs with that expected by integrating the accretion history of seed black holes \\citep{marconi04, shankar04}. Finally, popular semi-analytic models of galaxy formation and evolution \\citep{kh00, marulli08}, coupled to hydrodynamical simulations of galaxy mergers \\citep{hop06}, propose that nuclear activity is triggered during major mergers of gas-rich galaxies and that at its early stage, the AGN is embedded within optically thick gas shrouds. Despite these theoretical progresses, the cosmological evolution and luminosity function of CT objects is currently unknown, and in the synthesis models of the XRB it has been usually assumed to be equal to that of less obscured objects. Whether this is the case or not can only be determined by obtaining statistically significant samples of distant CT objects and by comparing them with the local samples (see e.g. \\citealt{c04} and \\citealt{rdc08} for reviews on nearby, {\\it bona-fide} CT objects).  The number of techniques devised to select CT AGN is rapidly growing, following the technological development of efficient detectors across the electromagnetic spectrum. These diverse selection techniques have allowed the first estimates of the space density of CT AGN in different redshift and luminosity regimes.  Very hard X-ray selection, i.e. at energies above 10 keV, is unaffected by absorption up to a few $\\times 10^{24}$ cm$^{-2}$, but, because of the still limited instrumental sensitivity, is mainly sampling the nearby Universe. The population of CT AGN detected by INTEGRAL/IBIS \\citep{tueller08} and Swift/BAT \\citep{malizia09} at $z<0.02$ is indeed producing only a tiny fraction ($<$1\\%) of the cosmic XRB.  Deep X-ray surveys in the more accessible 2-10 keV band are revealing large populations of heavily obscured objects in the redshift range $z\\sim[0.5-4]$. The generally low photon statistics of the detected sources, however, prevent an accurate spectral analysis and $N_H$ measurement. The CT nature of a faint X-ray source is therefore often inferred from the characteristics of the reprocessed spectrum, like the presence of a prominent fluorescence $K\\alpha$ iron line over a flat continuum. Examples of distant CT AGN selected with this technique have been found by \\citet{tozzi06} and \\citet{geo07,geo09} in the \\chandra\\ Deep Fields. The number of CT AGN candidates detected in the deep X-ray surveys appears to be in rather good agreement with the XRB models predictions.  A very recent and promising approach to select CT candidates in the distant Universe is based on their strong mid-IR flux, where most of the absorbed radiation should be re-emitted \\citep{alejo05}. Recently, \\citet{daddi07}, \\citet{fiore08, fiore09} and \\citet{alex08} located heavily obscured AGN in objects showing 24$\\mu$m emission in excess of that expected from dust heated by stellar processes. By stacking the \\chandra\\ data of these mid-IR-excess objects, a very hard X-ray spectrum was observed, reminiscent of CT obscuration. These studies span a broad AGN luminosity range ($L_X\\sim10^{42-45}$ erg s$^{-1}$), but mostly sample populations of objects at $z\\sim 2$. The measured space density of CT AGN at these high redshifts is in general as large as expected from XRB synthesis models or possibly even larger (see eg. \\citealt {trei09_ct}).  Another way to select obscured QSOs is through their high-ionization narrow optical emission lines, which are thought to be produced on physical scales (from $\\sim0.1$ to a few kpc) mostly free from nuclear obscuration. Recently, the [O IV]26$\\mu m$ line has been used to select obscured AGN among galaxies observed with {\\it Spitzer}/IRS. However, since this line quickly moves out of the observable IR bands as redshift increases, this selection mostly concerns the nearby Universe \\citep{diamond09, rigby09}. The most commonly used marker of obscured nuclear activity therefore remains the [O~III]5007 emission line, which is strong, falls in the optical domain, and allows object selection up to $z\\sim 0.8$.  The 2-10 keV to [O~III]5007 flux ratio (X/OIII) has been often used as a diagnostic for heavy obscuration in sources with poor X-ray photon statistics \\citep{maio98, cappi06, panessa06}, being low X/OIII ratios ($\\lesssim 3$, see e.g. Fig.~4 of \\citealt{cappi06}) highly suggestive of heavy nuclear absorption.  Based on the [O~III]5007 emission line, \\citet{zak03} and \\citet{reyes08} identified in the Sloan Digital Sky Survey (SDSS) a population of obscured QSOs at a median redshift of $z\\sim0.3$, at least as abundant as that of type-1 QSOs at the same redshifts \\citep{reyes08}. These results have been extended to lower luminosities by \\citet{bongiorno09}, who measured the luminosity function of [O III]-selected type-2 AGN in the zCOSMOS spectroscopic survey \\citep{lilly07}, finding that the fraction of obscured AGN is decreasing with luminosity, in agreement with what is observed in X-ray surveys.  X-ray observations of small samples drawn from the \\citet{zak03} catalog, suggest that about half of luminous type-2 QSOs (log$L_{O III}>9.3\\;L_{\\odot}$) could be CT \\citep[hereafter V10; see also \\citealt{lamastra09} for X-ray observations of lower luminosity SDSS type-2 AGN]{ptak06, v06, v10}. Since [O III] selection is likely missing objects in which also the Narrow Line Region (NLR) is extincted (like e.g. in the prototype CT AGN NGC~4945 and NGC~6240), the estimated type-2 QSO abundances should be considered as lower limits \\citep[V10]{reyes08}.  In this work we explore the possibility of using the high-ionization [Ne~V]3426 emission line, rather than the [O~III]5007 line, as a tracer of obscured nuclear activity. Despite being on average a factor of $\\sim 9$ weaker than [O~III]5007 \\citep{fo86, zak03} and suffering stronger dust extinction, the [Ne~V]3426 line is commonly observed in nearby Seyfert galaxies and, given that high energy photons ($\\gtrsim 0.1$ keV) are required to further ionize NeIV, it is considered an unambiguous sign of nuclear activity \\citep[e.g.][]{schmidt98}. In addition, the [Ne~V]3426 emission line is observable up to $z\\sim 1.5$ before being redshifted out of the optical bands, whereas the [O~III]5007 line is observable only up to $z\\sim 0.7-0.8$. Indeed, only 13 out of 887 [O~III]-selected QSOs in the \\citet{reyes08} sample lie at $z>0.7$, with only 3 at $z>0.8$. [Ne~V]-selection may then be used to reveal nuclear activity in obscured sources at $z\\sim 1$, i.e. at the epoch where most of the XRB light is thought to be produced.  The structure of the paper is the following: in Section 2 we present and discuss the sample of nearby Seyfert galaxies used to calibrate the relation between [Ne~V] and X-ray emission (the details of the sample are given in the Appendix). In Section 3 we present the X/NeV diagnostic diagram and apply it to obscured and unobscured QSO population drawn from the SDSS. In Section 4 we present \\chandra\\ observations of a sample of 9 [Ne~V]-selected obscured QSOs at $z\\sim 1$ in the SDSS and use the X/NeV diagnostic to estimate the fraction of CT objects among them. In Section 5 we discuss efficiency and biases of [Ne~V] selection together with its application to sky areas with deep optical spectroscopy and X-ray coverage. In the same Section, the evidence of enhanced star formation in obscured QSOs at z=0.4-1.5 is also highlighted. The conclusions are drawn in Section 6.  \\begin{figure*}[t] \\begin{center} \\includegraphics[width=12cm]{4039f1.ps} \\caption{Observed 2-10 keV to [Ne~V]3426 luminosity ratio (X/NeV) vs absorption column density for a sample of 74 Seyfert galaxies in the local Universe. The solid line shows the expected trend obtained by starting from the mean X/NeV ratio $\\langle$X/NeV$\\rangle$ observed in unobscured objects (i.e. those plotted at log$N_H=20$) and progressively obscuring the X-ray emission with increasing $N_H$ (up to log$N_H=25.5$) as plotted on the y-axis. Lower limits at log$N_H=24$ refer to Compton-Thick objects observed only below 10 keV, while datapoints at log$N_H>24$ and lower limits at log$N_H=25$ refer to CT objects observed also above 10 keV, for which a more stringent determination of the column density is possible. The darker (lighter) shaded region is obtained by making the same computation as above but starting at $\\pm1\\sigma$ ($\\pm 90\\%$) around $ \\langle$X/NeV$\\rangle$ (see text for details). The region at low ($<15$) X/NeV values is essentially populated only by CT objects.} \\label{xnev_local} \\end{center} \\end{figure*} ",
        "conclusions": "We have presented a diagnostic diagram to identify heavily obscured, Compton-Thick AGN candidates at $z\\sim 1$ based on the ratio between the 2-10 keV flux and the [Ne~V]3426 emission line flux (X/NeV). The diagnostic was calibrated on a sample of 74 local Seyfert galaxies and then applied to populations of type-1 and type-2 QSOs at different redshifts (from $z\\sim 0.1$ to $z=1.5$) selected from the SDSS. The main results obtained in this work can be summarized as follows. \\\\  \\noindent $\\bullet$ The observed X/NeV ratio is found to decrease with increasing absorption: the mean X/NeV ratio for unobscured Seyferts is about 400, about 80\\% of local Seyferts with X/NeV$<100$ are obscured by column densities above $10^{23}$\\cm\\, and essentially all objects with observed X/NeV $<15$ are Compton-Thick. \\\\  \\noindent $\\bullet$ We considered a sample of 83 blue type-1 QSOs and 21 [O III]-selected type-2 QSOs in the SDSS which have been observed in the X-rays and show significant [Ne V] detection. It was verified that they follow the same X/NeV vs X-ray absorption trend which is observed for local Seyferts. Furthermore, SDSS type-2 QSOs classified either as Compton-Thick or Compton-Thin on the basis of their X/OIII ratio, would have been mostly classified in the same way based on the X/NeV ratio. \\\\  \\noindent $\\bullet$ The X/NeV diagnostic was used to investigate the obscuration of 9 SDSS obscured QSOs in the redshift range $z=[0.85-1.31]$, which is not accessible through [O III] selection. The 9 objects were selected by means of their prominent [Ne~V]3426 line ($EW>4$\\AA), and \\chandra\\ snapshot observations for 8 of them were obtained (one object is from the archive). Based on the X/NeV ratio, complemented by X-ray spectral analysis, only 2 objects appear good Compton-Thick QSO candidates. However, when considering the 4 genuine narrow-line objects only (FWHM of the MgII line $\\lesssim 2000$ km $s^{-1}$), the efficiency in selecting Compton-Thick QSOs through the [Ne~V] line is about 50\\% (2/4), which is more similar, despite the large uncertainties, to what is achieved with [O~III] selection (60-70\\%; \\citealt{v10}). \\\\  \\noindent $\\bullet$ We verified that neither extinction nor anisotropy corrections on the [Ne~V] emission would affect our conclusions and that the X/NeV diagnostic is therefore a good method to identify clean, despite not complete, samples of heavily obscured AGN. We discussed the possibility of applying the X/NeV diagnostic to objects in sky areas with deep optical spectroscopy and X-ray coverage. This will allow to identify Compton-Thick Seyferts at $z\\sim 1$, i.e. those objects which are thought to be responsible for a large fraction of the ``missing\" X-ray background. \\\\  \\noindent $\\bullet$ Finally, the optical emission line properties of [Ne~V]-selected QSOs were compared with those of other SDSS populations of obscured and unobscured QSOs. By restricting the analysis to objects in the same redshift (and luminosity) range $z$=[0.4-1.5], we found evidence that the ratio between the [O~II]3727 and [Ne~V]3426 luminosity increases with obscuration. This correlation is interpreted as evidence of enhanced star formation in obscured QSOs, which is consistent with current popular scenarios of BH-galaxy coevolution.  {"
    },
    "1005/1005.1929_arXiv.txt": {
        "abstract": "Precise measurement of the angular power spectrum of the Cosmic Microwave Background (CMB) temperature and polarization anisotropy can tightly constrain many cosmological models and parameters. However, accurate measurements can only be realized in practice provided all major systematic effects have been taken into account. Beam asymmetry, coupled with the scan strategy, is a major source of systematic error in scanning CMB experiments such as Planck, the focus of our current interest. We envision Monte Carlo methods to rigorously study and account for the systematic effect of beams in CMB analysis. Toward that goal, we have developed a fast pixel space convolution method that can simulate sky maps observed by a scanning instrument, taking into account real beam shapes and scan strategy. The essence is to pre-compute the ``effective beams'' using a computer code, ``Fast Effective Beam Convolution in Pixel space'' ({\\tt FEBeCoP}), that we have developed for the Planck mission. The code computes effective beams given the focal plane beam characteristics of the Planck instrument and the full history of actual satellite pointing, and performs very fast convolution of sky signals using the effective beams. In this paper, we describe the algorithm and the computational scheme that has been implemented. We also outline a few applications of the effective beams in the precision analysis of Planck data, for characterizing the CMB anisotropy and for detecting and measuring properties of point sources. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.3092_arXiv.txt": {
        "abstract": "{Type Ia supernovae (SNe Ia) play a key role in measuring cosmological parameters, in which the Phillips relation is adopted. However, the origin of the relation is still unclear. Several parameters are suggested, e.g. the relative content of carbon to oxygen (C/O) and the central density of the white dwarf (WD) at ignition. These parameters are mainly determined by the WD's initial mass and its cooling time, respectively. Using the progenitor model developed by Meng \\& Yang, we present the distributions of the initial WD mass and the cooling time. We do not find any correlation between these parameters. However, we notice that the range of the WD's mass decreases, while its average value increases with the cooling time. These results could provide a constraint when simulating the SN Ia explosion, i.e. the WDs with a high C/O ratio usually have a lower central density at ignition, while those having the highest central density at ignition generally have a lower C/O ratio. The cooling time is mainly determined by the evolutionary age of secondaries, and the scatter of the cooling time decreases with the evolutionary age. Our results may indicate that WDs with a long cooling time have more uniform properties than those with a short cooling time, which may be helpful to explain why SNe Ia in elliptical galaxies have a more uniform maximum luminosity than those in spiral galaxies. ",
        "introduction": "% \\label{sect:1} As one of the most widely used distance indicators, type Ia supernovae (SNe Ia) show their importance in determining cosmological parameters, which resulted in the discovery of the accelerating expansion of the universe (\\citealt{RIE98}; \\citealt{PER99}). The result was exciting and suggested the presence of dark energy. At present, SNe Ia are proposed to be cosmological probes for testing the evolution of the dark energy equation of state with time and testing the evolutionary history of the universe (\\citealt{RIESS07}; \\citealt{KUZNETSOVA08}; \\citealt{HOWEL09}). They were even chosen to check the consistency of general relativity (\\citealt{ZHAOGB10}). When SNe Ia are applied as a distance indicator, the Phillips relation is adopted, which is a linear relation between the absolute magnitude of SNe Ia at maximum light and the magnitude drop of the B light curve during the first 15 days following the maximum (\\citealt{PHI93}). This relation implies that the brightness of SNe Ia is mainly determined by one parameter. It is generally agreed that the amount of $^{\\rm 56}$Ni formed during the supernova explosion dominates the maximum luminosity of SNe Ia (\\citealt{ARN82}), but the origin of the different amount of $^{\\rm 56}$Ni for different SNe Ia is still unclear (\\citealt{POD08}). Some numerical and synthetical results showed that metallicity has an effect on the final amount of $^{\\rm 56}$Ni, and thus the maximum luminosity (\\citealt{TIM03}; \\citealt{TRA05}; \\citealt{POD06}; \\citealt{BRAVO10}) and there do be some observational evidence of the correlation between the properties of SNe Ia and metallicity (\\citealt{BB93}; \\citealt{HAM96}; \\citealt{WAN97}; \\citealt{CAP97}; \\citealt{SHA02}). However, the metallicity seems not to have the ability to interpret the scatter of the maximum luminosity of SNe Ia (\\citealt{TIM03}; \\citealt{GALLAGHER08}; \\citealt{HOWEL09b}). \\citet{NOM99, NOM03} suggested that the ratio of carbon to oxygen (C/O) of a white dwarf at the moment of explosion is the dominant parameter for the Phillips relation. The higher the C/O, the larger the amount of nickel-56, and then the higher the maximum luminosity of SNe Ia. The C/O ratio is a function of the initial mass of the WD, which then is related to the progenitor system of the SNe Ia. By comparing theory and observations, the results of \\citet{MENGXC09} and \\citet{MENGXC10a} upheld this suggestion. \\citet{LESAFFRE06} carried out a systematic study of the sensitivity of ignition conditions for H-rich Chandra single degenerate exploders on various properties of the progenitors, and suggested that the central density of the WD at ignition may be the origin of the Phillips relation (see also \\citealt{POD08}). These authors noticed that the more massive and/or the cooler the CO WD is when accretion begins, the higher the central density is at ignition. The central density is then also related to the progenitor system. When one simulates the explosion of SNe Ia, the C/O and the central density are always set to be free parameters (\\citealt{ROPKE06}). It is thus interesting to analyze whether there is a correlation between the C/O and the central density. In addition, analyzing how the initial mass of the CO WD and its cooling time vary with the delay time is also an interesting task (\\citealt{GREGGIO10}). The purpose of this paper is to check these interesting problems.  In section \\ref{sect:2}, we describe our model. We show the results in section \\ref{sect:3} and give discussions and conclusions in sections \\ref{sect:4}. ",
        "conclusions": "\\label{sect:4} In this paper, we do not find a correlation between the initial WD mass and its cooling time. Since the C/O ratio is a function of the initial WD mass and the central density for a WD with given initial mass is mainly determined by the cooling time (\\citealt{NOM99, NOM03}; \\citealt{LESAFFRE06}), our results may imply that the C/O ratio and the central density at ignition are free parameters when simulating an SNe Ia explosion. However, our results still provide a constraint when simulating these types of explosions, i.e. the cooling time of WDs with an initial mass less than 1 $M_{\\odot}$ is generally shorter than 1 Gyr, but it may be as long as 15 Gyr for WDs with an initial mass larger than 1 $M_{\\odot}$ (see Fig. \\ref{mwdcoolt}). Because a high initial WD mass means a lower C/O, and a massive WD and a long cooling time leads to a high central density at ignition (\\citealt{NOM99, NOM03}; \\citealt{LESAFFRE06}), the above result could imply that WDs with a high C/O usually have a lower central density at ignition, but those having the highest central density at ignition generally have a lower C/O.  We also checked the effect of metallicity on the distributions of the initial WD mass and the cooling time by $Z=0.001$ (\\citealt{MENGXC10}) and no significant effect was found. Theory and observations did confirm that the effect of metallicity cannot explain the scatter in the maximum luminosity of SNe Ia (\\citealt{TIM03}; \\citealt{GALLAGHER08}; \\citealt{HOWEL09b}). Metallicity should, at most, be the secondary parameter for the Phillips relation, i.e. it is the origin for the scatter of the Phillips relation (\\citealt{POD06, POD08}). Furthermore, even regarding the secondary parameter, there is not a consensus. \\citet{MAZZZALI06} suggested that a variation of the relative content of ($^{\\rm 54}$Fe+$^{\\rm 58}$Ni) versus $^{\\rm 56}$Ni may be responsible for the observed scatter of the Phillips relation. However, \\citet{KASEN09} argued that the breaking of spherical symmetry is a critical factor in determining both the Phillips relation and the observed scatter around it. Then, the origin of both the Phillips relation and the scatter of the relation still should be investigated carefully.  We found that the range of the initial WD mass decreases and the average WD mass increases with the cooling time, which is similar to the relation between the initial WD mass and the delay time found by \\citet{MENGXC10a}. In addition, the scatter of the cooling time also decreases with delay time. These results may indicate that the difference among the WDs with s short cooling time could be huge, but the properties of WDs with a long cooling time might be more uniform (see from Fig. \\ref{mwdcoolt} that all the WDs with a cooling time longer than several Gyr have a mass larger than 1.0 $M_{\\odot}$). Since a long cooling time is equivalent to a long delay time (see Fig. \\ref{coolage}), our results may imply that the properties of SNe Ia with long delay times might be more uniform than those with short delay times. It is widely known that there exists a scatter of the maximum luminosity of SNe Ia, and the scatter is affected by its environment. The most luminous SNe Ia always occur in spiral galaxies, but both spiral and elliptical galaxies are hosts for dimmer SNe Ia, which lead to a dimmer mean peak brightness in elliptical than in spiral galaxies, i.e. the maximum luminosity of SNe Ia hosted in elliptical galaxies is more uniform (\\citealt{HAM96}). These properties may be qualitatively interpreted by the model developed by \\citet{MENGXC10a}, at least if the C/O ratio is the origin of the Phillips relation. However, the quantitative study by \\citet{ROPKE06} showed that the C/O has only little impact on the amount of produced $^{\\rm 56}$Ni. It should be noticed that the results in \\citet{ROPKE06} are model dependent, and are much different from those of \\citet{NOM99, NOM03}. Considering that the central density is mainly determined by the initial mass of the WD and its cooling time, the results in Figs. \\ref{mwdcoolt} and \\ref{coolage}, i.e. the average values of the WD mass and the cooling time increase with the delay time, which means that the average of the central density increases with the delay times of SNe Ia. A high central density at ignition leads to a larger amount of the produced $^{\\rm 56}$Ni (\\citealt{ROPKE06}). The effect of the central density on the amount of $^{\\rm 56}$Ni seems then to be opposite with observations of \\citet{HAM96}. Furthermore, the variation of the amount of $^{\\rm 56}$Ni derived from the central density only amounts to about 7\\%, which can not be used to interpret the variation in the maximum luminosity of SNe Ia (\\citealt{ROPKE06}). Perhaps, the C/O ratio, the central density and the metallicity all contribute to the variation of the maximum luminosity of SNe Ia (\\citealt{ROPKE06b}). Then, which is the dominant parameter for the Phillips relation is still an open queation. In addition, as the cooling/delay time increases, the CO WDs become more and more degenerate, and even crystallization may occur (\\citealt{FONTAINE01}). What is the effect of the crystallization on the amount of $^{\\rm 56}$Ni should be an interesting problem.  \\normalem"
    },
    "1005/1005.3747_arXiv.txt": {
        "abstract": "{The blazar \\source has been previously detected by the VERITAS and MAGIC telescopes in the very high energies. The new detection of VERITAS from December 2008 to April 2009 proves that \\source is not static, but shows short-time variability.}{We show that the time variability may be explained in the context of a self-consistent synchrotron-self Compton model, while the long time observation do not necessarily require a time-resolved treatment.}{The kinetic equations for electrons and photons in a plasma blob are solved numerically including Fermi acceleration for electrons as well as synchrotron radiation and Compton scattering.}{The light curve observed by VERITAS can be reproduced in our model by assuming a changing level of electron injection  compared to the constant state of \\source. The multiwavelength behaviour during an outburst becomes comprehensible by the model.}{The long time measurements of VERITAS are still explainable via a constant emission in the SSC context, but the short outbursts each require a time-resolved treatment.} ",
        "introduction": "Blazars are a special class of active galactic nuclei (AGN) exhibiting a spectral energy distribution (SED) that is strongly dominated by nonthermal emission across a wide range of wavelengths, from radio waves to gamma rays, and rapid, large-amplitude variability. The source of this emission is presumably the relativistic jet emitted at a narrow angle to the line of sight to the observer.\\\\In high-peaked BL Lac objects (HBLs) the SED shows a double hump structure as the most notable feature with the first hump in the UV- to X-ray regime and the second hump in the gamma-ray regime. Indeed, a substantial fraction of the known nearby HBLs have already been discovered with Cherenkov telescopes like H.E.S.S., MAGIC or VERITAS. The origin of the first hump is mostly undisputed: nonthermal, relativistic electrons in the jet are emitting synchrotron radiation. The origin of the second hump is still controversially debated. Up to now two kinds of models are discussed: leptonic \\citep[e.g.][]{maraschi92} and hadronic \\citep[e.g.][]{mannheim93} ones, which are mostly applied for other subclasses of blazars.\\\\Another important feature of AGNs in general and HBLs in particular is their strong variability. The dynamical timescale may range from minutes to years. This requires complex models, which obviously have to include time dependence, but this gives us also the chance to understand the mechanisms that drive AGNs. We will apply a self-consistent leptonic model to new data observed for the source \\source, because those are the ones favoured for HBLs.\\\\\\\\The source HBL \\source has been discovered as a candidate BL Lac object on the basis of its X-ray emission and has been identified with the X-ray source \\mbox{2A 1219+30.5} \\citep{wilson79,ledden81}. For the first time, \\source  has been observed at VHE energies using the MAGIC telescope in January 2005 \\citep{magic1218} and later from VERITAS \\citep{veritas1218}. Coverage of the optical/X-ray regime is provided by BeppoSAX \\citep{beppo05} and SWIFT \\citep{swift07}, unfortunately the data are not always simultaneous. During the observations from December 2008 to April 2009 VERITAS also observed \\source showing a time-variability \\citep{veritas2010}. The observations from the MAGIC telescope have previously been modelled by \\citet{michl1218}. \\citet{veritas2010} claim that their new observations exhibiting variability challenge the previous models. We will show that a timedependent model using a self-consistent treatment of electron acceleration is able to model the new VERITAS data.\\\\\\\\We present the kinetic equation, which we solve numerically, describing the synchrotron-self Compton emission (Sect. \\ref{sec:model}). In Sect.~\\ref{sec:results} we apply our code to \\source, taking the VERITAS data into account and give a set of physical parameters for the most acceptable fit. Finally, we discuss our results in the light of particle acceleration theory and the multiwavelength features. ",
        "conclusions": "\\label{sec:discussion} Our results clearly show that the latest observations from the VERITAS telescope for \\source still agree with a constant (steady state) emission from a SSC model when averaged over a long observation period. This is due to the relatively moderate variability of \\source compared to the observation time.\\\\The variability may be well explained in the context of the self-consistent treatment of acceleration of electrons in the jet. We are aware that an outburst of the timescale of roughly five days as measured from \\source does not necessarily require a shock in jet model, which scales down to a few minutes depending on the SSC parameters \\citep{weidinger2010b}, but may also be explained as e.g. different accretion states. Nevertheless the fundamental statement remains the same: long time observation of slightly variable blazars will result in a steady state emission, while an average over a single outburst will, of course, result in a significantly different SED for the source. We are not yet able to rule out different emission models or even complex geometries of the emitting region. But we are able to model the influence of short outbursts of a source on the SED and the lightcurves in the different energy bands selfconsistently.\\\\ The VERITAS collaboration only shows an integrated spectrum for \\source, which is due to the low flux of the source and the photon index behaviour of the combined high-states. This integrated spectrum does not show strong variations with regard to the known low-state observed by MAGIC. Our model now predicts a clear change in the spectrum, which is indicated by the dashed line in Fig. 1, which shows the average over one outburst with a slight, currently not detectable spectral softening in the VHE range, while the synchrotron peak in the BeppoSAX/SWIFT regime remains spectrally constant. This situation changes for shorter and/or stronger outbursts of an overall timescale of hours, which will result in measurable spectral evolutions in all energy regimes when considered with the presented model. Furthermore the time-resolved SEDs during a flare are comprehensible with our model. Hence with better time-resolved spectra or/and better multiwavelength coverage it should be possible to prove this model, and if the model is indeed applicable it will be a good tool to investigate the whole SED during an outburst without having all energy regimes observationally covered.\\\\\\textit"
    },
    "1005/1005.5482_arXiv.txt": {
        "abstract": "The global-scale interior magnetic field $\\BI$ needed to account for the Sun's observed differential rotation can be effective only if confined below the \\cz\\ in all latitudes including, most critically, the polar caps. Axisymmetric solutions are obtained to the nonlinear magnetohydrodynamic equations showing that such polar confinement can be brought about by a very weak downwelling flow $U\\sim10^{-5}$\\cms\\ over each pole.  Such downwelling is consistent with the \\hsc\\ evidence. All three components of the magnetic field $\\B$ decay exponentially with altitude across a thin, laminar ``magnetic \\cl'' located at the bottom of the \\tc\\ and permeated by spiralling field lines. With realistic parameter values, the thickness of the \\cl\\ $\\sim10^{-3}$ of the Sun's radius, the thickness scale being the magnetic advection--diffusion scale $\\CLthick=\\eta/U$ where the magnetic (ohmic) diffusivity $\\eta\\approx 4.1\\times10^{2}$\\cmms. Alongside baroclinic effects and stable thermal stratification, the solutions take into account the stable compositional stratification of the \\HSL, if present as in today's Sun, and the small diffusivity of helium through hydrogen, $\\chi\\approx0.9\\times10^{1}$\\cmms. The small value of $\\chi$ relative to $\\eta$ produces a double boundary-layer structure in which a ``\\hsl'' of smaller vertical scale $(\\chi/\\eta)^{1/2}\\CLthick$ is sandwiched between the top of the \\HSL\\ and the rest of the \\cl. Solutions are obtained using both \\san\\ and purely numerical, finite-difference techniques. The \\chl\\ flows are magnetostrophic to excellent approximation. More precisely, the principal force balances are between Lorentz, Coriolis, pressure-gradient and buoyancy forces, with relative accelerations negligible to excellent approximation. Viscous forces are also negligible, even in the \\hsl\\ where shears are greatest. This is despite the kinematic viscosity being somewhat greater than $\\chi$. \\ We discuss how the \\cl s at each pole might fit into a global dynamical picture of the solar \\tc.  That picture, in turn, suggests a new insight into the early Sun and into the longstanding enigma of solar lithium depletion. ",
        "introduction": "\\label{sec:intro}  This paper analyses a new family of laminar magnetostrophic flows that may be important for confining the interior magnetic field $\\BI$ needed to explain the Sun's differential rotation. As illustrated in \\fig\\ref{fig:helioseismology}, the differential rotation observed within the \\cz\\ goes over into near-\\solid\\ rotation within the radiative, stably stratified interior, via a thin shear layer called the ``\\tc'' much of which is also stably stratified. The need for the interior field $\\BI$ has been argued elsewhere \\citep[][hereafter GM98]{McIntyre94, RudigerKitchatinov97, GM98}; the main arguments are briefly recalled below. The observational evidence together with many ideas about the \\tc\\ are reviewed and further referenced in a recent major compendium, \\emph{The Solar Tachocline} \\citep{Hughesetal07}, and further discussed in the second edition of Mestel's \\emph{Stellar Magnetism} \\citep{Mestel11}.  By a ``confined'' $\\BI$ we mean a field most if not all of whose lines are contained beneath the \\cz, and held there against magnetic (ohmic) diffusion. Such confinement is well known to be necessary in order for the field to help enforce \\solid\\ rotation in the interior \\citep[\\eg][]{Ferraro37,MestelWeiss87, CharbonneauMacGregor93, MacGregorCharbonneau99}, and thereby keep the \\tc\\ thin. Confinement against magnetic diffusion requires fluid motion. So, besides magnetic effects, a realistic theory of confinement must take account of Coriolis effects, stable stratification, baroclinicity, and thermal relaxation. Without these effects we cannot correctly describe, for instance, the overall torque balance, which necessarily involves mean meridional circulations (MMCs) as well as Maxwell stresses.    The first attempt at a tachocline theory was that of \\citet{SpiegelZahn92}. It included all the above effects except $\\BI$. \\ \\citet{RudigerKitchatinov97} included $\\BI$ but omitted the other effects.  The first attempt to include all of them was that of GM98, in a line of investigation further developed by \\citet{GaraudGaraud08}. Meanwhile, the dynamical importance of compositional as well as thermal stratification \\citep[\\eg][]{Mestel53} was suggested for tachocline theories \\citep{McIntyre07}. In particular, the helium settling layer beneath the tachocline is nearly impermeable to MMCs because of the small diffusivity of helium through hydrogen. This near-impermeability of compositionally stratified regions has been called the ``\\muchoke'' \\citep{MestelMoss86}. The reality of the Sun's \\HSL\\ is strongly indicated both by standard solar-evolution models and by helioseismology \\citep[\\eg][]{ JCD-etal93, Ciacio-etal97, ElliottGough99, JCDThompson07}.\\footnote{ Also Christensen-Dalsgaard \\& Gough 2011, in preparation. } As will be seen, this combination of circumstances gives rise to some new and interesting fluid dynamics.  The need for the interior field $\\BI$ arises from a well known difficulty with non-magnetic theories. They tend to spread the strong differential rotation of the convection zone down into the radiative interior. Although sometimes disputed, this is a robust and well-understood consequence of thermal relaxation, interacting with Coriolis effects and gyroscopically-pumped MMCs \\citep{Haynes-etal91,SpiegelZahn92,Elliott97,McIntyre07,GaraudBrummell08}. As shown by \\citeauthor{SpiegelZahn92} and confirmed by \\citeauthor{Elliott97}, this downward spreading or burrowing would have produced a \\tc\\ far thicker than observed. The accompanying MMC, acting throughout the Sun's lifetime, would also have prevented the \\HSL\\ from forming.  To counter the burrowing tendency and to allow the interior to rotate \\solid ly, angular momentum has to be transported somehow from the low-latitude tachocline to the high-latitude tachocline.  The non-magnetic horizontal eddy viscosity proposed for this purpose by \\citeauthor{SpiegelZahn92} is inconsistent with the properties of non-magnetic stratified turbulence known from many studies of the terrestrial atmosphere \\citep[][\\& refs.]{McIntyre94,McIntyre03}. Angular-momentum transport by internal gravity waves is a physically possible alternative \\citep[\\eg][\\& refs.]{Schatzman93,Zahn-etal97, RogersGlatzmaier06,CharbonnelTalon07}. However, it is highly improbable as the main mechanism because, by itself, it has no natural tendency to produce \\solid\\ rotation at all latitudes and depths \\citep[\\eg][]{PlumbMcEwan78}.  A suitably-shaped magnetic field can, by contrast, naturally produce the required angular momentum transport, via the Alfv\\'enic elasticity of the field lines. A suitable shape is one in which the field lines link low latitudes to high latitudes within the \\tc. The simplest such shape --- simplest by virtue of its axisymmetry --- is that suggested schematically in \\fig\\ref{fig:sphere}, in which the linkage is via a time-averaged field whose lines thread the \\tc, forming the superficial part of a global-scale interior dipole stabilized by a deep toroidal field \\citep[\\eg][]{BraithwaiteSpruit04}. Such an interior dipole has a diffusive lifetime somewhat greater than the Sun's lifetime of around $4.5\\times10^9$\\yr. The dipole imposes an Alfv\\'enic ``Ferraro constraint'' on the interior.  It is this constraint that helps to enforce the interior's \\solid\\ rotation \\citep[\\eg][]{Ferraro37,MestelWeiss87, MacGregorCharbonneau99}.   \\begin{figure} \\centering \\subfigure[][]{ \\label{fig:helioseismology} \\psfrag{cz}{\\begin{tabular}{c}convection\\\\[-0.1em]zone\\end{tabular}} \\psfrag{rz}{\\begin{tabular}{c}radiative\\\\[-0.1em]interior\\end{tabular}} \\psfrag{tc}{tachocline} \\includegraphics[clip=true,bb=118 -12 614 510,height=6cm]{Schou-etal-bw.eps} } \\hspace{1cm} \\subfigure[][]{ \\label{fig:sphere} \\includegraphics[clip=true,bb=14 -152 878 1048,height=6cm]{spherebw.eps} } \\caption{\\footnotesize (a) The Sun's differential rotation deduced from helioseismic data using inverse methods \\citep[adapted from][]{Schou-etal98}. The radiative interior rotates approximately \\solid ly with angular velocity $\\OmegaI = 2.7\\times10^{-6}$\\sm, or 435\\,\\nHz. Within the \\cz, the angular velocity increases with colatitude through 350, 400, 450\\,\\nHz\\ (heavy contours) to a maximum just under 470\\,\\nHz\\ at the equator. \\newline (b) Schematic illustration showing the top of the radiative interior (inner sphere) and the time-averaged magnetic field threading the \\tc\\ just above. The cutaway outer sphere indicates the top of the \\tc, whose depth has been exaggerated. Poloidal magnetic field lines emerge from the interior in high latitudes and are wound up into their curved shapes by the \\tc's differential rotation, acting against turbulent eddy diffusion. A prograde torque is transmitted from low to high latitudes along these field lines. The slow polar and fast equatorial rotation are indicated by the darker shadings of the outer sphere. The dashed lines indicate the latitudes at which the rotation of the \\cz\\ matches that of the interior. } \\end{figure}   The field lines shown in \\fig\\ref{fig:sphere} emerge from the interior (light-grey sphere) near the north pole and, after threading their way through the \\tc, re-enter the interior near the south pole. They return northward through an interior ``apple-core'' region, not shown, surrounding the rotation axis. It is crucial that the field lines emerging from the interior bend over toward the horizontal as they enter the \\tc. They must be prevented from extending upward through the polar cap, as occurs when magnetic diffusion dominates \\citep[\\eg][]{BraithwaiteSpruit04,BrunZahn06}. The curved shapes of the field lines in \\fig\\ref{fig:sphere} are evidently such as to transport angular momentum from low to high latitudes by means of persistent Alfv\\'enic torques, exactly as required to prevent the \\tc's MMCs from burrowing into the interior and thickening the \\tc.   The time averaging envisaged in \\fig\\ref{fig:sphere} conceals a plethora of fast processes, including the 22-year dynamo cycle, convective overshoot, and other turbulent processes arising from various instabilities in the \\tc. All these are fast relative to the timescales on which the mean structure of the \\tc\\ is maintained, $\\sim10^5$\\yr\\ or more. We presume that the fast processes have two important consequences.  The first is to produce a turbulent magnetic diffusivity that stops the field lines being wound up arbitrarily tightly by the shear in the \\tc, keeping the curved shapes shown.\\footnote{Of course the persistent angular momentum transport from the curved $\\B$ lines could be supplemented by equally persistent contributions coming directly from MHD-turbulent stresses \\citep[\\eg][]{Spruit02,GilmanCally07,ParfreyMenou07}. }   The second important consequence is that, away from the poles, the field lines are held down, and held approximately horizontal, by turbulent ``\\mfp'' from the convective overshoot layer. The effectiveness of such flux pumping can be strongly argued from several lines of evidence, including three-dimensional direct numerical simulations, with varying emphasis on the role of turbulent anisotropy and of vertical gradients of density and turbulent intensity \\citep[\\eg][\\& refs.]{Tobias-etal01, KitchatinovRudiger06} \\citep[see also \\S3 of][for a historical review]{Weiss-etal04}.  Near the poles it is less clear that \\mfp\\ will be effective in confining the field. At least its effectiveness for near-vertical magnetic fields has not, to our knowledge, been convincingly demonstrated. However, as argued for instance in GM98, there are good reasons in any case (\\S\\ref{sec:downwelling} below) to expect the \\tc's MMC near the poles to take the form of weak but persistent downwelling. This suggests that the field can, in any case, be confined in the polar caps through an advection--diffusion balance, the kind of balance argued for heuristically in GM98. The purpose of this paper is to show in detail, by solving an appropriate set of nonlinear magnetohydrodynamic equations, that such polar confinement by downwelling is indeed possible in a physically realistic model, applicable to the Sun both today and early in its main-sequence lifetime.    A large family of axisymmetric nonlinear solutions showing polar confinement has been obtained using two different techniques. The first technique is \\san\\ in a sense to be explained, and the second is numerical on a \\mbox{2-dimensional} grid. The solutions are to be regarded as candidate solutions for possible flows in the real Sun, all showing confinement in the sense that the total magnetic field strength $|\\B|$ dies off exponentially with altitude, thanks to downward advection acting against upward diffusion. In this sense the poloidal and toroidal field components are both well confined. We call these flows ``\\cl s''. They are not to be confused with the \\tc\\ itself. Rather, they occupy relatively thin regions at the bottom of the \\tc\\ and are much more weakly sheared, with relatively long timescales $\\sim 10^5$\\yr.  The detailed dynamics involves not only magnetic advection, stretching, twisting and diffusion but also a near-perfect balance between Lorentz, Coriolis, pressure-gradient and buoyancy forces (\\S\\S\\ref{sec:model-equations}ff.). Thus the \\chl\\ flows are magnetostrophic, like certain flows that have been studied in connection with models of the Earth's liquid core \\citep[\\eg][\\& refs.]{Kleeorinetal97}, though different in most other respects. For instance the latter flows are viscous but unstratified: buoyancy forces and thermal diffusion are absent. In the \\chl\\ flows studied here, by contrast, viscosity turns out to be wholly unimportant while buoyancy and thermal diffusion are crucial, along with magnetic diffusion.  \\begin{figure} \\centering \\psfrag{U}[l][]{$\\mathbf{u}$} \\psfrag{B}[tl][B]{$\\mathbf{B}$} \\psfrag{W}[][t]{$_{\\phantom{\\ii}}\\OmegaI$} \\psfrag{TC}[][]{\\sc \\tc} \\psfrag{cl}[][]{\\begin{tabular}{c}confinement\\\\[-0.1em]layer\\end{tabular}} \\psfrag{sl}[][]{\\begin{tabular}{c}helium\\\\[-0.1em]sublayer\\end{tabular}} \\psfrag{HSL}[][]{\\sc\\begin{tabular}{c}helium\\\\[-0.1em]settling\\\\[-0.1em]layer\\end{tabular}} \\includegraphics[clip=true,bb=-205 150 818 675,width=10cm]{new-section.eps} \\caption{\\footnotesize The magnetic confinement layer near the north pole in a model for today's Sun. The field strength $|\\B|$ falls off exponentially with altitude $z$. \\ The toroidal components of $\\B$ and velocity $\\bu$ are not shown. The streamlines with arrows show the downwelling responsible for the confinement. If the downwelling were switched off, the field near the pole would diffuse and become nearly vertical, as illustrated for instance in \\citet{BrunZahn06}. Compositional stratification is indicated by shading. The plot is from a numerical solution; the corresponding \\san\\ solution looks almost identical. The horizontal and vertical axes are colatitudinal distance $r$ and altitude $z$ in units of $\\CLthick$, the advection--diffusion scale defined in (\\ref{eq:CLthickdef}). With typical parameter values, the scale $\\CLthick$ is of the order of a fraction of a megametre, $\\sim10^{-3}$ of the Sun's radius. } \\label{fig:section} \\end{figure}     \\Fig\\ref{fig:section} gives a preview of a typical \\chl\\ flow, seen in vertical cross-section. It shows the poloidal velocity and magnetic field components from a numerical solution. The emerging magnetic field lines are bent over within the \\cl, as required to fit into the global picture sketched in \\fig\\ref{fig:sphere}. The magnetic field $\\B$ has a toroidal component, not shown in the figure, imparting spiral shapes to the three-dimensional field lines and providing the prograde Alfv\\'enic torque demanded by the global picture, in balance with a retrograde Coriolis torque on the equatorward flow.  The vertical and colatitudinal distances in \\fig\\ref{fig:section} are shown in units of the magnetic advection--diffusion scale \\vspace{-0.15cm} \\begin{equation} \\CLthick \\definedas \\eta/U \\,, \\label{eq:CLthickdef} \\end{equation} say, where $U$ is the magnitude of the downwelling and $\\eta$ is the magnetic diffusivity. Throughout this paper, we assume that the \\chl\\ flow is laminar, and therefore use molecular or microscopic diffusivity values (\\S\\ref{sec:model-equations}). Issues of stability or instability lie beyond the scope of this paper but, close to the pole at least, there appears to be a strong case for stability, to be argued in a future paper, arising from the smallness of the scale $\\CLthick$. Under reasonable assumptions, $\\CLthick$ is only a fraction of a megametre, far smaller than the thickness of the overlying \\tc\\ which latter, by contrast, is probably unstable and indeed turbulent, as already mentioned \\citep[\\eg][]{Spruit02,GilmanCally07,ParfreyMenou07}.   In the present-day Sun's \\HSL, the top of which corresponds to the shaded region in \\fig\\ref{fig:section}, a downward gradient of helium concentration reinforces the stable stratification due to the sub-adiabatic temperature gradient. Because the diffusivity of helium through hydrogen, $\\chi\\approx0.9\\times10^{1}$\\cmms, is much less than the magnetic diffusivity $\\eta\\approx4.1\\times10^2$\\cmms, the \\HSL\\ is nearly impermeable to the \\chl\\ flow. Helium advection and diffusion are comparable only in the extremely thin ``\\hsl'' marked in \\fig\\ref{fig:section}. In this and other respects, all the solutions in the present paper supersede those described in a first report on this work \\citep[][hereafter WM07]{WM07}. For instance, in WM07 we took $\\chi$ to be zero, implying a \\hsl\\ of vanishing thickness. We also took $\\nu$, the kinematic viscosity, to be zero and allowed a finite slip velocity at the top of the \\HSL, assuming that this slip velocity would in reality be resolved into a weak Ekman layer. However, the solutions presented here show that, on the contrary, no Ekman layer forms. The slip discontinuity is replaced by a smooth velocity profile across the \\hsl\\ and, as will be shown in \\S\\ref{sec:helium}, the flow stays essentially inviscid.            The plan of the paper is as follows. In \\S\\ref{sec:downwelling} we summarize the reasons for expecting persistent downwelling over the poles. In \\S\\ref{sec:model-equations} we present the model equations and in \\S\\ref{sec:self-similar} the \\san\\ solutions.  Those solutions rely on assuming a self-similar horizontal structure that is asymptotically valid in the limit of strong stable stratification. The same limit was taken in WM07.    The validity of the strong-stratification limit is assessed in \\S\\S\\ref{sec:CL-scalings} and~\\ref{sec:helium}, which take a thorough look at the dynamical balances and scalings in the \\cl\\ and \\hsl\\ respectively. Strong stable stratification means that the thermal and compositional stratification surfaces are ``flat'', meaning gravitationally horizontal, to sufficient approximation in some region surrounding the poles, which for reasonable parameter values can be quite large in horizontal extent, up to tens of degrees of colatitude. Within the \\hsl, the low magnetic Reynolds number and flat geometry cause the momentum balance to take on the character of flow in a porous medium, as fluid pushes horizontally past the field lines.  As already indicated, true viscous effects are negligible everywhere, even in the sublayer.  Boundary conditions for the numerical solutions are discussed in \\S\\ref{sec:boundary-conditions}. The numerical solutions themselves are presented and discussed in \\S\\ref{sec:num-solutions}. They provide cross-checks with the \\san\\ solutions plus additional insights. In particular, they directly demonstrate the flatness of the stratification surfaces by solving the full equations, for finite stratification. The solutions allow the stratification surfaces to tilt as they may, but confirm that the departures from flatness are indeed small when the stratification is realistically strong. In \\fig\\ref{fig:section}, for instance, the departures from flatness are barely visible.  In \\S\\ref{sec:comparison} we discuss a subtlety that arises when comparing the \\san\\ and numerical solutions in the upper part of the flow. The dynamical balances aloft become delicate as the Lorentz and Coriolis forces become vanishingly small. The effects of truncation error and other small effects thus complicate the comparison. However, this is something of an academic point because of our expectation that, in reality, the \\chl\\ solutions will need to be matched to a turbulent \\tc\\ aloft, a task that remains a challenge for the future.   In \\S\\ref{sec:noHSL} we show that the presence of the \\HSL\\ is not crucial to our \\chl\\ model. The interior field $\\BI$ is sufficient by itself to turn the flow equatorward, and the field remains confined in much the same way. That result has relevance to the Sun's early main-sequence evolution. It explains for instance how the burrowing tendency could have been held in check from the start, allowing the \\HSL\\ to form. In the concluding discussion, \\S\\ref{sec:conclusions}, we consider the implications for early solar evolution and lithium depletion. ",
        "conclusions": "\\label{sec:conclusions}   We cannot yet claim to have a complete \\tc\\ theory.  Indeed, the \\cl\\ and \\hsl\\ form only two pieces of a complicated jigsaw puzzle. Other aspects of that jigsaw include the way in which the \\cl\\ matches upward to the relatively large negative shear in the bulk of the \\tc, and the way in which the baroclinic temperature anomalies induced by the \\tc's MMC fit into the perturbed global-scale heat flow. In particular, without putting the whole jigsaw together we cannot quantitatively predict the thickness of the \\tc. Nor can we predict the precise shapes of the vertical profiles of $\\B$ and $\\bu$ in the \\cl. Those profile shapes depend on matching to conditions not only aloft but also equatorward, where stratification surfaces and field lines extend into colatitudes outside the polar downwelling regions.  However, the results obtained here give us the first fully consistent model of polar field confinement, as such, together with insight into how it could work in today's Sun. The results cover a wide range of possible downwelling values and interior field strengths (end of \\S\\ref{sec:CL-scalings}). We have also shown, in \\S\\ref{sec:noHSL}, how confinement could have worked in the early Sun. The dynamics is similar apart from the slightly deeper penetration of the MMC in the absence of the \\HSL\\ and \\hsl. We can use the resulting insights, alongside our well-established understanding of the gyroscopic pumping of MMCs, to say something new about the early Sun and the solar lithium-burning problem.  Standard solar-evolution models predict surface lithium abundances higher than observed by a factor $\\sim10^2$ \\citep[\\eg][]{Vauclair-etal78}. The reason is that the standard models mix material down to the bottom of the \\cz\\ but no further.  To destroy lithium, material from the \\cz\\ must be mixed or circulated to somewhat greater depths and therefore to somewhat higher temperatures, beyond those at the bottom of today's \\tc. However, there is no evidence of depletion of the \\cz's beryllium, which is destroyed at only moderately higher temperatures than lithium. Further discussion and references may be found in \\citet{JCDGoughThompson92}, and in \\citet[][chap.\\ 6]{Wood10}. Here we argue that a quantitative version of the scenario sketched in \\fig\\ref{fig:sphere} has promise as a way of circulating material to the required depth in the early Sun, and no further, thus making sense of the high beryllium as well as the low lithium abundance.  As already mentioned in \\S\\ref{sec:intro}, the downwelling MMC in the polar \\tc\\ that makes field confinement possible can be regarded as due to a gyroscopically-pumped MMC trying to burrow downward, but held in check by its encounter with the interior field $\\BI$ and with the \\HSL, if present. If, in a \\te, we were to switch off the interior field $\\BI$, then the downwelling would spread or burrow to ever-increasing depths. The timescale for such burrowing is inversely proportional to $\\OmegaI^2$ \\citep[\\eg][\\eq(8.15)ff.]{McIntyre07}; one may think of rotational stiffness as strengthening the burrowing tendency.  Now, because the early Sun rotated much faster than today, not only would there have been no \\HSL\\ but, also, the burrowing tendency would have been much stronger than today, tending to push the bottom of the \\tc\\ downward. This reopens the possibility conjectured in {GM98} that there might have been a ventilated ``polar pit'' in which most of the \\cz's lithium, though not too much of its beryllium, was burnt during the first gigayear or two of the Sun's main-sequence evolution.  To take this further we again need to consider the way in which the \\cl\\ fits into the global picture.  It is arguable that the bottom of the entire polar downwelling region is depressed relative to its surroundings, forming not so much a ``pit'' as a shallow ``frying pan'', too shallow to burn lithium in today's Sun but possibly just deep enough in the early Sun.    \\begin{figure} \\centering \\includegraphics[width=7cm]{pit3.eps} \\caption{\\footnotesize Schematic drawing of the magnetic confinement layer and its immediate surroundings at the \\base\\ of the high-latitude \\tc. Close to the pole the interior magnetic field (solid lines) is confined by the downwelling MMC (dashed streamlines). The vertical scale has been greatly exaggerated.\\label{fig:surroundings}} \\end{figure} Here we need to distinguish the shape of the ventilated region from the shapes of the stratification surfaces, which latter must remain relatively flat, meaning close to the horizontal. \\Fig\\ref{fig:surroundings} sketches the way in which the \\cl\\ might fit into its surroundings near the bottom of the polar \\tc. The stratification surfaces are shown dotted. At the periphery of the polar downwelling region, the field lines (solid) spiral outward and upward from the \\cl\\ on their way to lower latitudes.  They will tend to splay out, as well as slanting upward, as they emerge from the downwelling region. The MMC will similarly slant upward, flowing approximately along the field lines (dashed streamlines).  This is because the splaying-out increases the magnetic Reynolds number beyond the order-unity values characteristic of the \\cl. Further out, the field lines must continue to rise through the \\tc\\ until they encounter the \\cz's overshoot layer, where they are held horizontal by turbulent \\mfp\\ as suggested in \\fig\\ref{fig:sphere}. On the way we must expect turbulent eddy fluxes to become increasingly important, decoupling the MMC's upwelling streamlines from the time-averaged field lines and leaving the upwelling free to spread over a wide range of latitudes, constrained only by mass conservation and global-scale heat flow.  Such a picture applies equally well to today's Sun and to the early Sun, the main difference being that the ventilated polar region (unshaded in \\fig\\ref{fig:surroundings}) is likely to have been pushed deeper in the early Sun with its much faster rotation, stronger burrowing tendency, and global-scale $|\\BI|$ values only modestly larger. The ventilated polar regions could well have been deeper by many tens of megametres, as required to burn lithium.\\footnote{This deepening is additional to the deepening of the \\cz\\ itself, in the early Sun relative to today's Sun, amounting to several more tens of megametres according to standard solar models \\citep[e.g.][]{Ciacio-etal97}. } The early Sun would have started to form its \\HSL\\ just below these ventilated, lithium-destroying polar regions, \\ie\\ just below the polar \\cl s. Then, with the gradual diminution of $\\OmegaI$ through solar-wind braking, the \\cl s, marking the bottom of the ventilated regions, would have retreated upward, and the top of the \\HSL\\ would have followed them upward as new helium strata formed.   Within the peripheral lightly-shaded region in \\fig\\ref{fig:surroundings}, into which the MMC does not penetrate, we suggest that ventilation is weak or nonexistent and that shear will be limited by the Ferraro constraint. The darker shading represents the top part of today's \\HSL.  As the lightly-shaded region expands upward and outward beyond the immediate surroundings sketched in \\fig\\ref{fig:surroundings}, through the \\tc\\ toward the overshoot layer, we may surmise that small-scale MHD instabilities will kick in \\citep[\\eg][\\& refs.]{Spruit02,GilmanCally07,ParfreyMenou07}, breaking the Ferraro constraint and blurring the distinction between the shaded and unshaded regions as turbulent eddy fluxes increase. So a larger-scale picture of the ``lithium frying pan'' would show its upward-sloping lower boundary becoming increasingly porous and indistinct at greater colatitudes.  The global \\tc\\ model that would be needed to test, and to begin to quantify, the foregoing speculations would have to describe  \\smallskip  \\begin{enumerate} \\item the precise way in which turbulent stresses in the convection zone and \\tc\\ gyroscopically pump the polar downwelling needed to confine $\\BI$ in polar latitudes; \\item the global-scale distribution of temperature and heat flow that fits in with the MMCs; \\item the turbulent \\mfp\\ by convective overshoot assumed to confine $\\BI$ in \\extrapolar\\ latitudes; \\item the extent to which the winding-up of the time-averaged toroidal field in \\extrapolar\\ latitudes (\\fig\\ref{fig:sphere}) is limited by turbulent eddy fluxes; \\item the reaction of the overlying turbulent layers to all of the above, especially the deficit in the \\cz's differential rotation governing the torques exerted from above, for instance via feedback on the strength of gyroscopic pumping of the MMC. \\end{enumerate}  \\smallskip  Progress on these formidable problems will depend on finding suitable ways to model the turbulent processes."
    },
    "1005/1005.5161_arXiv.txt": {
        "abstract": "We reconsider non-minimal $\\lambda\\,\\phi^4$ chaotic inflation which includes the gravitational coupling term $\\xi\\,\\mathcal{R}\\,\\phi^2$, where $\\phi$ denotes a gauge singlet inflaton field and $\\mathcal{R}$ is the Ricci scalar. For $\\xi \\gg 1$ we require, following recent discussions, that the energy scale $\\lambda^{1/4} m_P / \\sqrt{\\xi}$ for inflation should not exceed the effective UV cut-off scale $m_P / \\xi$, where $m_P$ denotes the reduced Planck scale. The predictions for the tensor to scalar ratio $r$ and the scalar spectral index $n_s$ are found to lie within the WMAP 1-$\\sigma$ bounds for $10^{-12} \\lesssim \\lambda \\lesssim 10^{-4}$ and $10^{-3} \\lesssim \\xi \\lesssim 10^2$. In contrast, the corresponding predictions of minimal $\\lambda\\,\\phi^4$ chaotic inflation lie outside the WMAP 2-$\\sigma$ bounds. We also find that $r \\gtrsim 0.002$, provided the scalar spectral index $n_s \\geq 0.96$. In estimating the lower bound on $r$ we take into account possible modifications due to quantum corrections of the tree level inflationary potential. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.3021_arXiv.txt": {
        "abstract": "The aim of this paper is to study the astrometric trajectory of microlensing events with an extended lens and/or source. We consider not only a dark lens but also a luminous lens as well. We find that the discontinuous finite-lens trajectories given by Takahashi will become continuous in the finite-source regime. The point lens (source) approximation alone gives an under (over)estimation of the astrometric signal when the size of the lens and source are not negligible. While the finiteness of the source is revealed when the lens transits the surface of the source, the finite-lens signal is most prominent when the lens is very close to the source. Astrometric microlensing towards the Galactic bulge, Small Magellanic Cloud and M31 are discussed, which indicate that the finite-lens effect is beyond the detection limit of current instruments. Nevertheless, it is possible to distinguish between self-lensing and halo lensing through a (non-)detection of the astrometric ellipse. We also consider the case where the lens is luminous itself, as has been observed where a lensing event was followed up with the \\textit{Hubble Space Telescope}. We show that the astrometric signal will be reduced in a luminous-lens scenario. The physical properties of the event, such as the lens-source flux ratio, the size of the lens and source nevertheless can be derived by fitting the astrometric trajectory. ",
        "introduction": "Most of the microlensing events detected to date are through photometric monitoring of the source flux. In this case, the information on the physical identity of the lens is reduced, because the only quantity one can retrieve from the light curve is the Einstein timescale $t_{\\mathrm{E}}$. $t_{\\mathrm{E}}$ is defined by the time required for the source to transit the angular Einstein radius $\\AERR$ of the lens \\citep{2000ApJ...542..785G}: \\begin{equation} t_{\\mathrm{E}} = \\frac{\\AERR}{|\\mbox{\\boldmath $\\mu_{rel}$}|},\\quad \\AERR = \\sqrt{k\\ML\\pi_{rel}},\\quad k\\equiv \\frac{4G}{c^2 \\mathrm{AU}}\\approx 8.14\\frac{\\mathrm{mas}}{M_{\\odot}}, \\label{eq.tE} \\end{equation} where \\mbox{\\boldmath $\\mu_{rel}$} is the relative lens-source proper motion, $\\ML$ is the mass of the lens, $\\pi_{rel}:= \\mathrm{AU}/(\\Dol^{-1}-\\Dos^{-1})$ is the relative lens-source parallax, $\\Dol$ and $\\Dos$ are distance to the lens and the source from the observer, respectively. Equation~(\\ref{eq.tE}) shows that the mass, distance, and velocity of the lens are degenerated into $t_{\\mathrm{E}}$.  To better constrain the lens properties, \\cite{1995A&A...294..287H}, \\cite{1995ApJ...453...37W} and \\cite{1995AJ....110.1427M} thus suggested to use astrometric microlensing. That is, to measure the centroid displacement of the two images during the course of microlensing. Former studies have shown that the trajectory of the centroid displacement will trace out an ellipse, and the size of the ellipse is proportional to the angular Einstein radius. Therefore, one can determine $\\AERR$ through the observation of such astrometric ellipses and constrain the relative lens-source proper motion. \\cite{1992ApJ...392..442G} has shown that if one can further measure the microlens parallax $\\pi_{\\mathrm{E}}=\\sqrt{\\pi_{rel}/(k\\ML)}$ form the light-curve distortion induced by the orbital motion of the Earth, the lens mass $\\ML$ and the relative lens-source parallax $\\pi_{rel}$ can be determined without ambiguity: \\begin{equation} \\ML = \\frac{\\AERR}{k\\pi_{\\mathrm{E}}},\\quad \\pi_{rel}=\\pi_{\\mathrm{E}}\\AERR. \\label{eq.M} \\end{equation}  The location of the lens can be derived as well if the distance to the source is well known, which is often the case towards the Galactic bulge and Magellanic Clouds.  The typical value of the astrometric microlensing signal for a source in the Galactic bulge and a 0.5 $M_{\\odot}$ lens located half-way to the source is of order of 0.1 mas, which is much larger than the astrometric accuracy of upcoming space missions such as \\textit{Space Interferometry Mission} \\citep[\\textit{SIM}; ][]{1997SPIE.2871..504A}, \\textit{Global Astrometric Interferometer for Astrophysics} \\citep[\\textit{GAIA}; ][]{1994SPIE.2200..599L} and ground-based instruments, e.g. \\textit{Phase Referenced Imaging and Micro-arcsecond Astrometry} \\citep[\\textit{PRIMA}; ][]{1998SPIE.3350..807Q}. \\textit{GAIA} will survey the whole sky with sources brighter than 20 mag in \\textit{V} band. It is expected to reach an astrometric accuracy of 30 $\\mu$as (150 $\\mu$as) with $V <$ 12 ($V <$ 16) for a single measurement \\citep{2002MNRAS.331..649B} and an estimated detection of $\\approx$ 1000 events \\citep{2000ApJ...534..213D}. Unlike \\textit{GAIA}, which only scans the sky with a pre-determined pattern, \\textit{SIM} can point to selectable targets and thus tracks the ongoing microlensing event upon request. The expected accuracy of \\textit{SIM} is 5 $\\mu$as (20 $\\mu$as) for $V <$ 12 ($V < $ 16) with 1-hour integration time \\citep{2008SPIE.7013E.151G}. While \\textit{SIM} and \\textit{GAIA} are scheduled to launch in the next few years, \\textit{PRIMA} has already been installed on the Very Large Telescope Interferometer (VLTI) and aims at achieving 10-$\\mu$as accuracy level in 30 min integration time provided a reference star within 10 arcsec and a 200-m baseline \\citep{2008NewAR..52..199D}.  In addition to the standard point-source point-lens (PSPL) microlensing, single-lens events revealing an extended source signal have also been observed photometrically \\citep[e.g.][]{1997ApJ...491..436A, 2004ApJ...617.1307J, 2004ApJ...603..139Y, 2006A&A...460..277C, 2009A&A...508..467B, 2009ApJ...703.2082Y, 2009arXiv0912.2312Z, 2010A&A...518A..51F}. \\cite{1998MNRAS.300.1041M} thus derived the astrometric trajectory of finite-source events with a point-lens (FSPL). On the other hand, \\cite{2003ApJ...595..418T} studied the centroid displacement of finite-lens effects but assuming a point-source (PSFL). Furthermore, \\cite{2002ApJ...579..430A} and \\cite{2009ApJ...695..200L} have investigated the combination of finite-source and finite-lens (FSFL) effects photometrically, but left aside the astrometric aspect. Since the FSFL light curve deviates from either the PSFL or FSPL, as shown by \\cite{2002ApJ...579..430A}, we are motivated to study the astrometric behaviour when both FS and FL effects are relevant. There are events where both the source and the lens are resolved by the \\textit{Hubble Space Telescope (HST)}. This implies that the lens can also be a star and implies a luminous-lens scenario rather than a dark lens \\citep[][]{2001Natur.414..617A, 2007ApJ...671..420K}, which is also the case for self-lensing. We thus consider the light contribution from the lens star and study the astrometric behaviour by allowing for a luminous lens.  This paper is organized as follows. In \\textsection~2 we introduce the theory of astrometric microlensing. We take into account the FS effects either with a uniform surface brightness source or with a more general surface brightness profile in \\textsection~3. We further include a dark lens with finite size in \\textsection~4. One might expect not only shadowing but also light contribution from the lens as well. Therefore, we allow for a luminous lens in \\textsection~5. The aforementioned properties of the microlensing system can be estimated by fitting the formula in \\textsection~5. A discussion of possible events with sources located in the Galactic bulge, Small Magellanic Cloud (SMC) and M31 is presented in \\textsection~6 followed by a summary in \\textsection~7. \\\\ ",
        "conclusions": "We have studied the astrometric aspects of microlensing by simultaneously including the FS and FL effects. Our results show that the astrometric signal is underestimated or overestimated by assuming PL or PS, respectively. While the FS effect is prominent when the lens transits the surface of the source, the FL effect is revealed when the lens is very close to the source, which would be in the self-lensing regime. In the context of the self-lensing scenario, where a background star is lensed by a foreground star, the light contribution from the lens is in general not negligible. We thus consider the luminous-lens scenario, which attenuates the signal of the centroidal displacement. Astrometric trajectories with a source located in the Galactic bulge, SMC, and M31 are discussed, which show that $\\AERR$ of halo-lensing events is at least one order of magnitude larger than that of self-lensing in SMC and M31. Our results also indicate that the finiteness of the lens is more likely to be revealed in the self-lensing scenario towards distant source located in Magellanic Clouds or M31, although it is very difficult to distinguish between PL and FL with current instruments."
    },
    "1005/1005.3810_arXiv.txt": {
        "abstract": "Dark energy is an important science driver of many upcoming large-scale surveys. With small, stable seeing and low thermal infrared background, Dome A, Antarctica, offers a unique opportunity for shedding light on fundamental questions about the universe. We show that a deep, high-resolution imaging survey of 10,000 square degrees in \\emph{ugrizyJH} bands can provide  competitive constraints on dark energy equation of state parameters using type Ia supernovae, baryon acoustic oscillations, and weak lensing techniques. Such a survey may be partially achieved with a coordinated effort of the Kunlun Dark Universe Survey Telescope (KDUST) in \\emph{yJH} bands over 5000--10,000 deg$^2$ and the Large Synoptic Survey Telescope in \\emph{ugrizy} bands over the same area. Moreover, the joint survey can take advantage of the high-resolution imaging at Dome A to further tighten the constraints on dark energy and to measure dark matter properties with strong lensing as well as galaxy--galaxy weak lensing. ",
        "introduction": "Antarctic Plateau, especially the Kunlun Station under construction by the Polar Research Institute of China (PRIC), provides a unique opportunity for wide field astronomical surveys targeting cosmological studies. Astronomical site survey of Dome A, Antarctica was enabled by the International Polar Year (IPY) endorsed PANDA program~\\citep{Yang09} led by the PRIC and the Chinese Center for Antarctica Astronomy (CCAA). Several international teams contributed to this effort. In particular, the power and on-site laboratory system built by the University of New South Wales (BUN'S) has provided the platform for all the site survey instruments.  In its two years' operation, the site survey effort proves practically all aspects of the theoretical expectations of the Dome A site for astronomical observations. Preliminary analyses show that the boundary layer of atmospheric turbulence to be around $10-20$ meters during the Antarctic winter~\\citep{ashley10}. Similar to the relatively better studied neighboring Dome C site~\\citep{fossat10}, the Dome A site may enjoy free atmospheric seeing conditions of about 0.3 arcsec seeing above this boundary layer, thus making it an ideal site for high angular resolution wide area surveys.  The temperature at Dome A is around $-60$ to $-70 \\degree{C}$, making it the coldest spot on the surface of the Earth. This implies a very low thermal background emission in the thermal infrared. The Dome A site is thus also the best site for astronomical observations in the near infrared wavelength region.  One other especially exciting property of the site is the lack of water vapors due to the high altitude and the low temperature. This implies that the site is also ideal for terahertz observations which have been impossible from any temperate sites on Earth.  While quantitative site properties are still under analyses and longer term monitoring are still needed to firmly establish the astronomical potential of the Dome A site, there is no doubt that a survey project at Dome A can be highly complimentary to programs such as the Large Synoptic Survey Telescope\\footnote{See \\url{http://www.lsst.org/}.} \\citep[LSST,][]{lsst09}, the Joint Dark Energy Mission\\footnote{See \\url{http://jdem.gsfc.nasa.gov/}.}, and Euclid\\footnote{See \\url{http://sci.esa.int/euclid/}.}. For example, a survey at Dome A can provide near infrared data that compliments the deep optical band survey of the LSST; a deep survey in the near infrared combined with the optical data from LSST can reveal high redshift objects at $z\\sim10$, which are not detectable in the optical. The Kunlun Dark Universe Survey Telescope (KDUST) is a 6-to-8-meter wide-area survey telescope being designed by the CCAA. The preliminary design includes a $3\\times3$ square degree optical camera with $0''.15$ pixel, and an infrared camera of $1\\times1$ square degree at $0''.1$/pixel optimized for $1-3.5$ $\\mu$m surveys.  One of the key science missions of KDUST is to investigate the mystery of the accelerated cosmic expansion \\citep{riess98,perlmutter99a} using multiple techniques. In this paper, we estimate how well an ideal 10,000 deg$^2$ \\emph{ugrizyJH} survey can constrain the dark energy equation of state (EOS) with weak lensing (WL), baryon acoustic oscillations (BAOs), and type Ia supernova (SNe) luminosity distances. These dark energy probes have different sensitivities to the cosmic expansion and structure growth as well as various systematic uncertainties in the observations, and hence are highly complementary to each other for constraining dark energy properties \\citep[e.g.,][]{knox05,zhan06d,zhan09}.  The SNe technique relies on the standardizable candle of the SNe intrinsic luminosity \\citep{phillips93} to measure the luminosity distance, $D_{\\rm L}(z)$. Dark energy properties can then be inferred from the distance--redshift relation. The BAO technique utilizes the standard ruler of the baryon imprint on the matter (and hence galaxy) power spectrum \\citep{peebles70, bond84} to measure the angular diameter distance, $D_{\\rm A}(z)$, and, if the redshifts are sufficiently accurate, the Hubble parameter, $H(z)$ \\citep{eisenstein98, cooray01b, blake03, hu03b, linder03, seo03}. The WL technique has the advantage that it can measure both $D_{\\rm A}(z)$ from the lensing kernel and the growth factor of the large-scale structure $G(z)$ \\citep{hu99,huterer02b,refregier03,takada04,knox06b,zhan09}.  Because of the excellent seeing condition and infrared accessibility at Dome A, KDUST has a number of advantages for the commonly used cosmological probes. For example, the signal-to-noise ratio for point sources is inversely proportional to the seeing. Thus, a 6 meter telescope at Dome A ($\\sim 0''.3$ median seeing in the optical) would be equivalent to a 14 meter telescope at a temperate site ($\\sim 0''.7$ seeing) for point-source observations with the same sky background level. An 8m KDUST could detect SNe out to redshift 3 in the $K_{\\rm dark}$ ($2.27$--$2.45\\mu$m, redward of $K$) band \\citep{kim2010}. Although high-z distances are not sensitive to conventional dark energy, they can be used to determine the mean curvature accurately, which, in turn, helps constrain dark energy EOS \\citep{linder05b,knox06c}. Moreover, even though dark energy is thought to be sub-dominant at high redshift, there is no direct evidence to prove one way or another. Measurements of SNe at $z > 2$ will provide crucial data for tests of early dark energy.  Small and stable seeing is particularly helpful for WL. One could resolve more galaxies at the same surface brightness limit, which reduces the shape noise for shear measurements. Fine resolution helps measure the shape accurately and reduce the shear measurement systematic errors. In addition, deep \\emph{JH} photometry can track the 4000\\AA{} break of an elliptical galaxy to $z \\sim 3$ and improve photometric redshifts (\\phz{}s) as well as systematic uncertainties in the \\phz{} error distribution \\citep{abdalla08}, which has a large impact on WL constraints on the dark energy EOS \\citep{huterer06,ma06,zhan06d}.  Adding \\emph{K} or $K_{\\rm dark}$ band will certainly improve galaxy \\phz{}s, especially at $z \\gtrsim 3$. However, currently planned multiband dark energy surveys use galaxies at $z \\lesssim 3$, and their concern is the confusion between $z\\lesssim 0.5$ ellipticals and $2 \\lesssim z \\lesssim 3.5$ star-forming galaxies, which is greatly mitigated by $u$ and $JH$ bands \\citep{abdalla08}. Another consideration is that Dome A is far more advantageous at $K_{\\rm dark}$ band than at $K$ band because of the low thermal background there, so $K_{\\rm dark}$ is likely to be chosen over $K$. This would leave a considerable gap in the wavelength coverage and reduce the already-small gain on \\phz{}s in the useful redshift range for dark energy investigations. Therefore, we do not discuss utilities of wavebands beyond $H$ in this paper. Nevertheless, the $K_{\\rm dark}$ band is crucial for a broad range of other sciences and will be an important aspect of the KDUST survey.  This paper is organized as follows. Section~\\ref{sec:survey} discusses the survey plan of KDUST including that of its pathfinder in context of the LSST survey. We then consider a joint KDUST and LSST survey in \\autoref{sec:de} for constraining the dark energy EOS with BAO, WL, and type Ia SN techniques. The results are presented in both two-parameter space where the dark energy EOS is parameterized as $w(z)=w_0+w_az(1+z)^{-1}$ and in model-independent principle component space. The conclusion is drawn in \\autoref{sec:con}. ",
        "conclusions": "\\label{sec:con} Dome A offers a very competitive site for studying dark energy. Given the amount of resources required to build a large telescope and run a massive survey there, one must give the highest priority to programs that cannot be easily carried out elsewhere. Thus, a reasonable strategy is to focus on NIR imaging and collaborate with other surveys for optical data. Using LSST as an example, we show that a high-resolution 5000--10,000 deg$^2$ KDUST survey in \\emph{yJH} bands could improve LSST BAO+WL constraints on the dark energy EOS parameters $w_0$ and $w_a$ by reducing the \\phz{} and shear measurement systematics. A SNAP-like SN sample plus a large local and nearby SN sample from KDUST would further boost the DETF FOM by more than a factor of two.  In addition to forecasts for the $w_0$--$w_a$ parametrization, we also apply a PCA approach to investigate the constraints on the dark energy EOS $w(z)$ in a model-independent way. We find that regarding the number of the constrained eigenmodes of $w(z)$, an ideal 10,000 deg$^2$ \\emph{ugrizyJH} survey, combined with \\emph{Planck}, can constrain 7 eigenmodes, while KDUST+LSST can allow us to constrain 3 more modes.  We have not discussed dark energy probes such as strong lensing, cluster counting, and higher-order statistics of the same galaxy and shear data, which could further tighten the constraints on the dark energy EOS. Strong lensing constrains dark energy through the time delay effect as well as counting of strong lenses. It is also an excellent probe of dark matter halo structures and, hence, can be used to measure dark matter particle properties. With high-resolution imaging, one could extract more cosmological information from strong lensing observations. Therefore, Dome A could be particularly advantageous for strong lensing studies.  Dark energy forecasts depend crucially on the assumed properties of the survey data, including all the systematics. Dome A has many advantages over other ground sites and has an environment close to that in space. Hence, we use well-studied LSST and SNAP as references to make crude estimates of the data for this investigation. Further work and detailed modeling are needed to give a more realistic assessment of the Dome A site for studying dark energy."
    },
    "1005/1005.3217_arXiv.txt": {
        "abstract": "We analysed 866 observations of the neutron-star low-mass X-ray binary XTE J1701--462 during its 2006-2007 outburst. XTE J1701--462 is the only example so far of a source that during an outburst showed, beyond any doubt, spectral and timing characteristics both of the Z and atoll type. There are 707 RXTE observations ($\\sim 2.5$ Ms) of the source in the Z phase, and 159 in the atoll phase ($\\sim 0.5$ Ms). We found, respectively, pairs of kilohertz quasi-periodic oscillations (kHz QPOs) in 8 observations during the Z phase and single kHz QPO in 6 observations during the atoll phase. Using the shift-and-add technique we identified the QPO in the atoll phase as the lower kHz QPO. We found that the lower kHz QPO in the atoll phase has a significantly higher coherence and fractional rms amplitude than any of the kHz QPOs seen during the Z phase, and that in the same frequency range, atoll lower kHz QPOs show coherence and fractional rms amplitude, respectively, 2 and 3 times larger than the Z kHz QPOs. Out of the 707 observations in the Z phase, there is no single observation in which the kHz QPOs have a coherence or rms amplitude similar to those seen when XTE J1701--462 was in the atoll phase, even though the total exposure time was about 5 times longer in the Z than in the atoll phase. Since it is observed in the same source, the difference in QPO coherence and rms amplitude between the Z and atoll phase cannot be due to neutron-star mass, magnetic field, spin, inclination of the accretion disk, etc. If the QPO frequency is a function of the radius in the accretion disk in which it is produced, our results suggest that in XTE J1701--462 the coherence and rms amplitude are not uniquely related to this radius. Here we argue that this difference is instead due to a change in the properties of the accretion flow around the neutron star. Regardless of the precise mechanism, our result shows that effects other than the geometry of space time around the neutron star have a strong influence on the coherence and rms amplitude of the kHz QPOs, and therefore the coherence and rms amplitude of the kHz QPOs cannot be simply used to deduce the existence of the innermost stable circular orbit around a neutron star. ",
        "introduction": "It is now more than 13 years ago that kilohertz (kHz) quasi-periodic oscillations (QPOs) were discovered (\\citealt{v1996}; \\citealt{st1996}) in neutron star (NS) low-mass X-ray binary (LMXB) systems. Interest in this phenomenon remains high because of the close correspondence between kHz QPO timescales and typical dynamical timescales of matter orbiting close to a NS. For this reason, kHz QPOs are potential tools to probe general relativity in the strong-gravitational field regime \\citep{v2005}, and constrain the NS equation of state \\citep{mil1998}.\\\\ Since the launch of the Rossi X-Ray Timing Explorer (RXTE) in 1995, kHz QPOs have been detected in about 30 NS LMXBs (for a review see \\citealt{v2005}). Most of these sources show two simultaneous kHz QPOs, usually called the lower and the upper kHz QPO, with frequencies that can drift as a function of time in the range 250-1200 Hz (\\citealt{v2004}). Studies of these kHz QPOs show that QPO frequencies are related to other properties of the source; e.g. on short time-scales (within a day or less) QPO frequencies are well correlated with the intensity of the source, whereas on long time-scales this correlation breaks down and intensity-frequency diagrams show the so-called ``parallel tracks'' (\\citealt{m1999}). The frequencies of the kHz QPOs correlate also with the position of the source in the colour-colour diagram, and with parameters of spectral components used to describe the X-ray spectrum of these sources  (\\citealt{w1997}, \\citealt{mv1999}, \\citealt{k1999}, \\citealt{ds2001}). Nevertheless, it is still unclear which physical parameters drive the QPO frequency, although there are indications that mass accretion rate, $\\dot{m}$, plays a key role \\citep*{mil1998}.\\\\ Several models have been proposed to explain the kHz QPOs (e.g., \\citealt{mlp1998}, \\citealt{sv1998}, \\citealt{a2003}), as well as the connection between high-frequency QPOs and other time variability usually present in power density spectra (for a review of variability at low frequencies see \\citealt{v2001}). Despite these efforts, there is still no single model that is able to explain in a self-consistent way all the QPO properties.\\\\ KHz QPOs are characterised by three parameters: centroid frequency $\\nu$, quality factor $Q=\\nu/FWHM$, where $FWHM$ is the full-width at half-maximum of the QPO, and fractional rms amplitude. Systematic analyses of these kHz QPO properties have been done for a large number of sources \\citep[e.g.][]{j2000,van2000,ds2001,mvf2001,h2002,van2002,dmv2003,b2005a,b2005b,a2005,m2006,bom2006}. Those studies show that, in each source the quality factor and the rms amplitude of the lower kHz QPO increase with the centroid frequency of the QPO until they reach a maximum value, after which they decrease as the frequency continues to increase (e.g. \\citealt*{mvf2001}, \\citealt*{dmv2003}, \\citealt{b2005b}; see \\citealt{m2006} for a compilation of results). The upper kHz QPO does not show the same trend; in this case the quality factor usually does not change with the centroid frequency while the rms amplitude stays more or less constant at lower frequencies and then decreases as the frequency increases (\\citealt{van2002}; \\citealt{van2003}; \\citealt{b2005a}; \\citealt{a2008}).\\\\ \\citet{b2005b} and \\citet{bom2006} interpreted the drop of the quality factor of the lower kHz QPO at high frequencies in the LMXBs 4U 1636--536 and 4U 1608--522 as a signature of the inner disk radius reaching the innermost stable circular orbit (ISCO), and starting from that assumption they estimated the mass and the radius of the compact object in these two systems. However, \\citet{m2006} argued against this idea and suggested that the drop of $Q$ and rms in individual sources might be related (at least in part) to changes of the properties of the accretion flow in these systems.\\\\ Following those results, here we investigate the properties of the kHz QPOs for the transient NS LMXB XTE J1701--462. This source was detected for the first time on 2006 January 18 with the All-Sky monitor on-board RXTE \\citep{r2006}. As reported by \\citet{h2007ate}, \\citet{l2009a}, \\citet{a2010} and \\citet{h2010}, this is the only source so far that showed both Z and atoll behaviour (for more details about the Z and atoll classes see \\citealt{hv1989}). The luminosity range covered by XTE J1701--462, from Eddington limit to quiescence, gives us a unique opportunity to study kHz QPO properties in different states and, more importantly, at different mass accretion rates in the same system, which could provide vital information to understand the origin and the mechanisms that drive the properties of these QPOs.\\\\ In section 2 we describe the observations and the data analysis, and in section 3 we present our results. In section 4 we discuss those results in the context of current ideas concerning the mechanisms behind the kHz QPOs in LMXBs, and in section 5 we summarise our conclusions. \\section[]{OBSERVATIONS AND DATA ANALYSIS} \\label{data} \\begin{table*} \\scalebox{0.87} \\centering \\begin{tabular}{c|||c|c|c|||c|c|c||} \\hline \\hline \\multicolumn{7}{c}{Z phase} \\\\ \\hline Obs ID & \\multicolumn{3}{c||}{$L_{\\ell}$} & \\multicolumn{3}{c||}{$L_{u}$}\\\\ \\hline & $Q_{\\ell}$ & rms$_{\\ell}\\%$ & $\\nu_{\\ell}$ (Hz) & $Q_{u}$ & rms$_{u}\\%$ & $\\nu_{u}$ (Hz) \\\\ \\hline 91442-01-07-09 & $35.3\\pm18.8$&$1.4\\pm0.2$&$642.7\\pm3.1$&$52.8\\pm38.8^{1}$&$<1.6$&$932.6\\pm7.9^{1}$ \\\\ \\hline 92405-01-01-02 & $22.6\\pm11.1$&$3.4\\pm0.5$&$615.1\\pm3.8$&$56.9\\pm28.8$&$2.9\\pm0.5$&$944.9\\pm3.7$\\\\ \\hline 92405-01-01-04 & $5.4\\pm2.8^{1}$&$< 3.5$&$502.4\\pm23.1^{1}$&$10\\pm3$&$3.1\\pm0.3$&$760.8\\pm6.4$ \\\\ \\hline 92405-01-02-03 & $11.9\\pm6.5$&$3.0\\pm0.5$&$623.9\\pm8.3$&$20.6\\pm11.1$&$2.7\\pm0.4$&$911.2\\pm8.4$\\\\ \\hline 92405-01-02-05 & $8.9\\pm6.9$ & $2.7\\pm0.4$& $595.9\\pm10.1$ & $6.2\\pm1.4$ & $4.1\\pm0.4$ & $850.5\\pm12.4$\\\\ \\hline 92405-01-03-05 &$8.5\\pm2.7$&$4.6\\pm0.5$&$612.5\\pm7.7$&$10.5\\pm3.5$&$4.6\\pm0.6$&$917.3\\pm10.6$\\\\ \\hline 92405-01-40-04 &$9.2\\pm3.2$&$2.8\\pm0.3$&$650.5\\pm7.1$&$11.1\\pm3.1$&$2.9\\pm0.3$&$911.0\\pm8.6$\\\\ \\hline 92405-01-40-05 &$8.1\\pm2.4$&$3.0\\pm0.3$&$637.9\\pm8.6$&$12.9\\pm3.5$&$3.0\\pm0.3$&$919.8\\pm7.6$\\\\ \\hline \\end{tabular} \\centering \\caption{Properties of the kHz QPOs detected in the Z phase of XTE J1701--462. Subscript letters $u$ and $\\ell$ denote lower and upper kHz QPOs, respectively. Parameters without errors represent 95\\% confidence level upper limits. All other errors reported represent 1$\\sigma$ confidence intervals.  $^{1}$These parameters are calculated for kHz QPOs with a significance level lower than 3$\\sigma$ (see the text for precise values).} \\label{tab2} \\end{table*} We analysed all the public data of the LMXB XTE J1701--462 collected with the Proportional Counter Array (PCA) on board of RXTE (\\citealt{b1993}; \\citealt{ja2006}). There are 866 observations of this source in the RXTE archive, for a total exposure time of $\\sim$ 3 Ms. During these observations the source showed several type-I X-ray bursts that we excluded from our analysis (see \\citealt{l2009b} for a detailed analysis of the bursts). \\subsection{Timing analysis} To search for kHz QPOs, we created Leahy-normalised power density spectra using Event mode data with 125$\\mu s$ time resolution covering the full PCA energy band, nominally from 2 to 60 keV. We created Fourier power density spectra from 16-seconds data segments, using 1/4096 s time resolution such that the frequency range is defined from 0.0625 Hz to 2048 Hz. We removed detector drop-outs; no dead-time correction or subtraction of background contribution were done to calculate the power density spectra. We created one averaged power density spectrum for each observation that we visually inspected to search for the presence of QPOs with characteristic frequencies in the range from 200 Hz to 1200 Hz. We found kHz QPOs in 14 out of the 866 observations that we analysed.\\\\ Following \\citet{a2010} and \\citet{h2010}, we considered that XTE J1701--462 was in the Z phase from its discovery in January 2006 \\citep{r2006} until the end of April 2007 (as reported by \\citealt{a2010} and \\citealt{h2010}, no clear boundary between Z and atoll phase has been found which makes this date just an approximation), when it started to behave as an atoll source until it went into a quiescence state. Using this division, there are 707 observations ($\\sim 2.5$ Ms) of XTE J1701--462 in the Z phase, and 159 observations ($\\sim 0.5$ Ms) in the atoll phase. KHz QPOs were detected in individual observations only in the horizontal branch during the Z phase, and in the lower banana in the atoll phase of the outburst \\citep[see][]{h2010}. From 14 observations with kHz QPOs, 8 belong to the Z phase (ObsIDs are reported in Table \\ref{tab2}) and 6 to the atoll phase (observations: 93703-01-02-04, 93703-01-02-11, 93703-01-02-05, 93703-01-02-08, 93703-01-03-00, and 93703-01-03-02).\\\\ A quick analysis of the observations showed clear differences in the properties of the QPOs between Z and atoll phases (see \\S3 for a detailed discussion). The QPOs were always weaker and broader in the Z than in the atoll phase. It is well known that the frequency of the kHz QPOs can change over tens of Hz in time intervals of a few hundred seconds \\citep[e.g][]{b1996}, and this can artificially broaden the QPO in the averaged power spectrum of long observations. Therefore, for each observation in which we found a kHz QPO, we divided the observation in smaller intervals to check if a significant QPO was still present, and whether the QPO frequency was changing. We found that in the Z phase it was not possible to detect a significant kHz QPO in power spectra of intervals shorter than a full observation: In all observations in the Z phase with kHz QPOs the QPOs were weak and broad over short time intervals. Therefore, for the rest of the analysis, for the Z phase we report QPO properties for the average power density spectrum of each observation. In the atoll phase, the kHz QPOs were significantly detected on time intervals shorter than a full observation, and in most cases the frequency was changing in time. Therefore, for the atoll phase, we decided to average power density spectra according to the frequency of the QPO. To do this we produced power density spectra of 16s of data, and averaged up to 15 of these power spectra to get a significant detection of the kHz QPO from which we measured the QPO frequency as a function of time. Finally we averaged power density spectra such that the frequency of the QPO was within a range of 10 Hz to 60 Hz (frequency intervals are reported in Table \\ref{tab1}). We shifted the frequency scale of all these power density spectra in order to align the frequencies of the QPOs to a constant value and then averaged the power density spectra to create one power density spectrum per selection (see \\citealt{m1998}). We fitted each Z and atoll average power density spectrum in the frequency range 200-1200 Hz using a constant to model the Poisson noise plus one Lorentzian to model the kHz QPO. It must be clarified that, for the atoll phase, the centroid frequencies of the QPOs reported in Table \\ref{tab1} are the mean frequencies within the interval selected, and the errors associated are the standard deviations of the selection.\\\\ In order to label the QPOs we used the standard convention introduced by \\citet{bpv2002} where each QPO is denoted with the letter \\emph{L} with a subscript that specifies the category. In this particular case we were interested in kHz QPOs so we used $L_{\\ell}$ and $L_{u}$ to identify the lower and the upper kHz QPO, respectively. Following this criterion all the characteristics associated with one QPO have the same label.\\\\ We estimated the fractional rms amplitude of QPOs in the atoll phase in different energy bands. For each observation with QPOs we first created power density spectra in 5 different energy intervals: 2-3 keV, 3-6 keV, 6-11 keV, 11-16 keV and 16-25 keV (we stopped at 25 keV because of the lack of sensitivity of the detector above that energy).  We shifted and added the power density spectra creating one single power density spectrum for each energy band and  we fitted these power density spectra as the atoll power density spectra previously described. To calculate the fractional rms amplitude we calculated the integral power of the Lorentzian and we renormalised it using the source and background count rate (see \\citealt{v1997}). \\\\ \\subsection{Spectral analysis} We calculated X-ray colours and intensity of the source using the Standard 2 mode data. We defined a hard colour as the count rate ratio in the energy bands 9.7-16.0 keV and 6.0-9.7 keV, and the intensity of the source as the count rate in the energy band 2.0-16.0 keV. To obtain the exact count rate in each of these bands we interpolated in channel space.\\\\ To correct for the gain changes and differences in the effective area between the proportional counter units (PCUs) as well as differences due to changes in the channel to energy conversion of the PCUs as a function of time, we normalised our colour and intensity by the Crab Nebula values obtained close in time to our observations and per PCU  (see \\citealt{k2004} and \\citealt{a2008} for details). Finally we averaged the normalised colours and intensities per PCU every 16 seconds using all available PCUs.\\\\ We also calculated the source luminosity for all 14 observations containing kHz QPOs. In order to do so, we first created a light curve for each observation, we then checked for X-ray bursts and detector drop-outs, and if any were present we eventually excluded them from the analysis. Using Standard 2 data we extracted energy spectra following the procedures described in the RXTE web page, and we added a 0.6\\% systematic error in quadrature to each channel. We fitted the energy spectra in Xspec in the energy range from 3 keV to 22 keV using a model consisting of a blackbody and a multi-colour blackbody for the Z observations, while we added also a broken power-law with the break energy fixed at 20 keV to fit the atoll observations \\citep[see][]{l2009a}. These model also include absorption from the interstellar medium toward the source and when necessary we added a Gaussian emission line at $\\sim$ 6.5 keV. The reduced $\\chi^{2}$ of our fits range from 0.6 to 1.1 (for 36 d.o.f.). From the best-fitting model we calculated the unabsorbed flux in the energy range 2-50 keV, setting the N$_{H}$ to zero. and creating an artificial response function for the full energy range 2 keV to 50 keV. We then estimated the luminosity assuming a distance $d= 8.8$ Kpc \\citep{l2009b} .\\\\ \\begin{table} \\centering \\begin{tabular}{c||c|c|c||c|} \\hline \\hline \\multicolumn{5}{c}{atoll phase}\\\\ \\hline Interval &Hz & $Q$& rms$\\%$ & $\\nu$(Hz)  \\\\ \\hline 1 & 600-660 & $69.9\\pm17.1$ & $9.8\\pm0.9$ & $640.9\\pm13.5$\\\\ \\hline 2 & 660-700 & $59.6\\pm15.6$ & $10\\pm1$ & $673.3\\pm14.5$\\\\ \\hline 3 & 700-750 & $98.2\\pm14.8$ & $9.2\\pm0.5$ & $716.8\\pm6.7$\\\\ \\hline 4 & 750-780 & $67.1\\pm19.2$ & $11.8\\pm1.2$ & $772.2\\pm4.2$\\\\ \\hline 5 & 780-800 & $106.9\\pm9.7$ & $9.9\\pm0.3$ & $793.9\\pm5.4$\\\\ \\hline 6 & 800-820 & $150.3\\pm20.9$ & $8.7\\pm0.4$ & $811.5\\pm6.6$\\\\ \\hline 7 & 820-830 & $114.9\\pm12.7$ & $9.3\\pm0.4$ & $827.3\\pm2.6$\\\\ \\hline 8 & 830-840 & $93.4\\pm8.7$ & $8.9\\pm0.3$ & $835.8\\pm2.6$\\\\ \\hline 9 & 840-850 & $100.6\\pm13.9$ & $7.9\\pm0.4$ & $846.2\\pm2.9$\\\\ \\hline 10 & 850-950 & $123.9\\pm36.8$ & $6.6\\pm0.7$ & $854.8\\pm4.3$\\\\ \\hline \\end{tabular} \\caption{Properties of the kHz QPOs detected in the atoll phase of XTE J1701--462. Column 2 shows the frequency selections used to create the intervals (see text for details). Column 3, 4 and 5 show the quality factor, the fractional rms amplitude and the frequencies of the kHz QPOs, respectively. All errors represent 1$\\sigma$ confidence intervals. } \\label{tab1} \\end{table} \\section[]{RESULTS} \\label{res} In Table \\ref{tab2} and Table \\ref{tab1} we report the quality factor $Q$, fractional rms amplitude and frequency of the kHz QPOs detected in all the RXTE observations available for XTE J1701--462, from January 2006 to August 2007.\\\\ In Table \\ref{tab2} we show the properties of the kHz QPOs during the Z phase. The significance level of these QPOs ranges between 3.2$\\sigma$ and 6$\\sigma$, except for $L_{u}$ in observation 91442-01-07-09 and $L_{\\ell}$ in observation 92405-01-01-04 which have, respectively, significances of 2.4$\\sigma$ and 2.3$\\sigma$. (The significances of kHz QPOs are given as the ratio of the integral of the power of the Lorentzian used to fit the QPO divided by the negative error of the power. As shown by \\citealt{bou2010}, this probably underestimates the true significance of the QPOs). For these two kHz QPOs we report upper limits. During the Z phase the lower and upper kHz QPOs show similar quality factors, on average around 15; also the fractional rms amplitude of the two QPOs is similar, between $\\sim$ 1.4\\% and 4.5\\%.\\\\ In Table \\ref{tab1} we report the properties of the 10 QPOs detected in the frequency-selected intervals in the atoll phase. These QPOs have significances between 5$\\sigma$ and 15$\\sigma$. The QPO frequency varies from 640 Hz to 860 Hz; the quality factor changes from about 60 up to 150. The fractional rms amplitude ranges from about 7\\% to 12\\%.\\\\ \\begin{figure} \\begin{center} \\includegraphics[width=80mm]{figure1.eps} \\caption{Left panel shows the frequency of the QPOs in the atoll phase of XTE J1701--462 as a function of the hard colour of the source. Right panel show the frequency of the QPOs as a function of the intensity of the source.} \\label{hard} \\end{center} \\end{figure} \\subsection{QPO identification} While for the Z phase it is easy to label the QPOs, for the atoll phase the identification is not straightforward since there we found just one QPO in each observation. A quick inspection of the values in Table \\ref{tab1} shows similar QPO properties between different intervals, which indicates that in all cases we are probably dealing with the same QPO. To progress further we compared $Q$ and rms values with those of other LMXBs where two simultaneous kHz QPOs have been studied (\\citealt{w1997}, \\citealt{w1998}, \\citealt{j1998}, \\citealt{b2005a}, \\citealt{b2005b}, \\citealt*{bom2006}); the coherence and rms amplitude are similar to those of the lower kHz QPO in other atoll sources.\\\\ According to \\citet{b2007} and \\citet{m1999}, lower and upper kHz QPOs in atoll sources follow two different tracks in the frequency-hardness diagram. Lower kHz QPOs are found when the source has low hard colour and the centroid frequency does not seem to be correlated with hard colour (as the frequency changes the hard colour changes slightly in a restricted range). On the contrary, upper kHz QPOs are found where the source has high hard colour that decrease as the QPO frequency increases \\citep[see Figure 3 in][]{b2007}. In Figure \\ref{hard} we plot the QPO frequency as a function of the hard colour and  intensity of the source. All points are concentrated within a narrow hard colour range, from 0.7 to 0.75, while the frequency ranges from 640 Hz to 860 Hz. The fact that the hard colour changes within a restricted range while the frequency moves in a range of about 200 Hz resembles what \\citet{b2007} show concerning the lower kHz QPO behaviour in the atoll source 4U 1636--53. This further supports our identification of the kHz QPOs in the atoll phase as lower kHz QPOs.\\\\ In order to understand the frequency-hardness diagram, we plot in the right panel of Figure \\ref{hard} the QPO frequency as a function of the source intensity. It is apparent from the graph that the data are divided into two different tracks characterised by intensity values which differ by a factor of $\\sim 2$. That trend resembles the so-called ``parallel tracks'' phenomenon which has been seen in many LMXBs \\citep[see e.g.][]{m1999}. Although the points in Figure \\ref{hard} are the result of frequency selection, it turns out that all the QPOs we combined to make the intervals from 1 to 4  and from 5 to 10 on Table \\ref{tab1} are concentrated in a time interval, respectively, of about 1 day and less than half a day. Moreover, those two groups are separated by more or less 4 days, which means the points in the right panel of Figure \\ref{hard} can in fact reflect the ``parallel tracks'' phenomenon.\\\\ \\begin{figure} \\begin{center} \\includegraphics[width=80mm]{figure2.eps} \\caption{Power density spectrum of all the atoll observation where QPOs are detected. This power spectrum is the result of the shift-and-add method. Two kHz QPOs are visible, the one at lower frequency (30 $\\sigma$ significant) was the one originally used ti shift and add the power spectra, while the one at higher frequencies (3.1$\\sigma$ significant, single trial) appears as result of the application of the method. } \\label{double} \\end{center} \\end{figure} We notice that QPOs at the highest intensities correspond to the ones with the highest frequencies in the frequency-hardness diagram, while QPOs at the lowest intensities correspond to those with the lowest frequencies in the frequency-hardness diagram. This can be interpreted as ``parallel tracks'' in the frequency-hardness diagram. A similar trend was seen in the LMXB 4U 1636--53 (see Figure 2 in \\citealt*{dmv2003}).\\\\ To search for a second (possibly weaker) kHz QPO in the atoll phase we apply the shift-and-add method \\citep{m1998}. As we described in section \\ref{data}, first we fit the centroid frequency of the kHz QPO in each individual atoll power density spectrum, then we shift all the power spectra such that the kHz QPO frequencies are aligned at the same value. Finally we average all the data into a single power density spectrum which we fit with a constant plus Lorentzian component for each QPO. This method corrects the frequency drift in time of the kHz QPOs, increasing their signal to noise ratio and their significance. If the second kHz QPO is at a (more or less) fixed distance from the first, this method adds all the data in a way that increases the significance of the second QPO, which may then become detectable.\\\\ The result of this procedure is presented in Figure \\ref{double}, which shows two simultaneous kHz QPOs, the one at lower frequency is the kHz QPO which we already knew, while the second one becomes detectable as a result of the shift-and-add method. The significance level of the second, upper, kHz QPO is 3.1$\\sigma$, which implies a marginal detection. Note however that the frequency difference between the kHz QPOs is $258\\pm13$ Hz, which is consistent with the value in the Z phase (see Table \\ref{tab2} and \\citealt{h2007}) and therefore there are no other trials involved in estimating this significance. As a result of this analysis, taking into account all the caveats, we find a pair of simultaneous kHz QPOs in the atoll phase of the XTE J1701--462 which strengthens our previous suggestion that the strong kHz QPOs detected in this phase are always the lower kHz QPO.\\\\ As described in section \\ref{data}, we calculate the fractional rms amplitude of the atoll QPOs in 5 different energy bands, 2-3 keV, 3-6 keV, 6-11 keV, 11-16 keV and 16-25 keV, respectively. We find fractional rms amplitudes of less than 10.9\\%, $6.9\\pm0.3$, $11.4\\pm0.2$, $17.5\\pm1.2$ and less than 21.3\\%, respectively (upper limits are given at 95\\% confidence). Those results, as already noticed in other sources (see Berger et al. 1996), show that the strength of the variability increases as the energy increases.\\\\ \\subsection{Amplitude and coherence of the kHz QPOs in the Z and atoll phases} Figure \\ref{Q_nu} shows the values of the QPO quality factor as a function of frequency. Black points represent measurements in the atoll phase while grey points are measurements in the Z phase. For the Z phase, empty and filled square symbols represent, respectively, lower and upper kHz QPOs. From the plot we notice a few interesting aspects: 1) QPOs in the atoll phase are on average 10 times more coherent than in the Z phase; 2) in the atoll phase $Q$ increases as the frequency increases, reaching a maximum value of about 150 at $\\sim$ 810 Hz; within the errors the behaviour is consistent with what has been observed in other sources \\citep[see][]{b2005a}. No such trend seems to be present in the Z phase; 3) comparing lower kHz QPOs in both phases it is clear that in the Z phase the lower kHz QPO appears at lower frequencies than in the atoll phase; 4) there is  an overlap in frequency around 640 Hz between Z and atoll QPOs with a significant mismatch in $Q$.\\\\ Figure \\ref{r_nu} shows the QPO rms fractional amplitude as a function of frequency. As in Figure \\ref{Q_nu}, black points denote measurements in the atoll phase, while grey points are QPOs in the Z phase. The rms amplitude of the lower kHz QPO in the atoll phase remains more or less constant around 10\\% as the QPO frequency increases from 640 Hz to 780 Hz and then drops rapidly at $\\sim$ 800 Hz. There is no evidence of a similar trend for the rms of the lower or upper kHz QPOs in the Z phase. As already noticed for the quality factor, the rms amplitude values show a clear difference between the atoll and Z phases: The lower kHz QPOs amplitude in the atoll phase is on average a factor of 2 higher than in the Z phase; this difference is also apparent in the region where atoll and Z QPOs overlap in frequency between 620 Hz and 660 Hz.\\\\ We test whether we may have missed any significant kHz QPO in our analysis. We calculate upper limits to the fractional rms amplitude in frequency ranges where we did not find kHz QPOs. To do so we take two power density spectra, one for each source phase, where no kHz QPOs were found. We fit those power density spectra using a model consisting of a constant to fit the Poissonian noise, and a Lorentzian to fit the QPO with fixed values for the centroid frequency and the quality factor $Q$. We fit the data and we estimate the upper limit of the amplitude of the Lorentzian using $\\Delta \\chi^{2}=2.7$, which corresponds to 95\\% confidence level. We repeat the same procedure shifting gradually the frequency of the Lorentzian to higher values until we cover the frequency range of interest.\\\\ For the analysis of atoll observation we use power density spectra of 256s, comparable to the time intervals over which we detected kHz QPOs in the atoll phase. We use two different quality factors, $Q=50$ and $Q=20$, which are comparable to the values we found for the kHz QPOs at low frequencies (see Table \\ref{tab1}). We find upper limits to the fractional rms amplitude that vary within the range 2.7\\%-5.2\\% and within the range 3\\%-7\\%, respectively for $Q=20$ and $Q=50$. From these values it appears that the rms amplitude of the lower kHz QPO in the atoll phase should decrease at frequencies lower than 640 Hz. This would be in agreement with the typical rms-frequency trend observed in most of the LMXBs (\\citealt*{mvf2001}; \\citealt{van2000, van2002, van2003}; \\citealt{ds2001}; \\citealt*{dmv2003}; \\citealt{b2005a}).\\\\ For the Z observation, we estimate the upper limits to the amplitude of a kHz QPO with typical atoll properties within short time intervals to check whether we could have missed such a QPO during the analysis. We calculate the upper limits in the frequency range 650-940 Hz assuming a kHz QPO with $Q=100$ and setting 3 values of the integration time, 128, 256 and 512 seconds. We find that the upper limits to the fractional rms amplitude vary between 1\\% and 4\\% for 128 seconds of integration time, between 0.8\\% and 4\\% for 256 seconds, and between 0.4\\% and 3.4\\% for 512 seconds. From this we conclude that if a kHz QPO with $Q=100$ and fractional rms amplitude similar to what we found for the atoll kHz QPOs was present in the Z observations, we would have detected it significantly in intervals as short as 128s.\\\\ ",
        "conclusions": "The transient source XTE J1701--462 is so far the only accreting neutron-star X-ray binary that changed from a Z into an atoll source during an outburst. Here we found that the properties of the kHz QPOs in XTE J1701--462 in the atoll and Z phases are significantly different: At approximately the same frequency (about 640 Hz), the coherence of the kHz QPOs is a factor of 2 larger, and the rms amplitude is a factor of 3 larger, in the atoll than in the Z phase (see Figure 3 and Figure 4, respectively). Furthermore, out of 707 observations in the Z phase, there is no single one in which the kHz QPOs have a coherence or rms amplitude similar to those seen when XTE J1701--462 was in the atoll phase, even though the total exposure time was about 5 times longer in the Z than in the atoll phase. If the kHz QPOs reflect the motion of matter in the accretion disk, and the QPO centroid frequency is related to the radius in the accretion disk where this motion takes place, our results show that, at least in XTE J1701--462 there is no unique relation between the radius of the disk at which the QPO is produced and the coherence and rms amplitude of the QPO.\\\\ It was already known that the kHz QPOs are broader and weaker in Z than in atoll sources (see M\\'endez 2006 for a compilation), but this is the first time that the same trend is observed in a single source. This result conclusively excludes things like neutron-star mass, magnetic field, spin, inclination of the accretion disk, etc., as the cause of this trend, since these parameters cannot change on time scales of one and half year. As we discuss below, the most likely reason for the difference of QPO coherence and rms amplitude between the Z and atoll phase in XTE J1701--462 is a change in the properties of the accretion flow around the neutron star where the QPOs are produced. Despite the the lack of information about the precise mechanism, our result shows that effects other than the geometry of space time around the neutron star have a strong effect on the QPO properties. If, as suggested by \\citet{b2005a, b2005c}, the ISCO affects the coherence and rms amplitude of the kHz QPOs, our result shows that there are other mechanisms that should also be taken into account to explain the trend seen in the data. For instance, in XTE J1701--462 the coherence and rms amplitude differ by  a factor of $\\sim 3$ to 10 between the Z and atoll phases similar, for instance, to the change of coherence and rms amplitude that has been proposed to be due to the ISCO in 4U 1636--53 \\citet{b2005b}.\\\\ \\citet{m2006} found that the the behaviour of the coherence and rms amplitude of the kHz QPOs in individual sources is similar to the behaviour of the maximum coherence and maximum rms of the kHz QPOs as a function of the luminosity in a sample of 12 NS-LMXBs. It is interesting to test to what extent this luminosity relation holds for the Z and atoll phases of XTE J1701--462, which are well separated in luminosity. In Figure \\ref{max} we combine the data points from the Z and atoll phases of XTE J1701--462 with those of the 12 NS-LMXBs studied by \\citet{m2006}.\\\\ \\begin{figure*} \\begin{center} \\includegraphics[angle=-90,width=120mm]{figure3.eps} \\caption{Quality factor $Q$ as a function of the QPO centroid frequency for XTE J1701--462. Black points represent QPOs in the atoll phase, and grey points represent QPOs in the Z phase. Empty and filled squares indicate lower and upper kHz QPOs, respectively. The two values of the coherence at frequencies $\\sim$500 Hz and $\\sim$930 Hz correspond to kHz QPOs with significances 2.3$\\sigma$ and 2.4$\\sigma$, respectively (see text).} \\label{Q_nu} \\end{center} \\end{figure*} \\begin{figure*} \\begin{center} \\includegraphics[angle=-90,width=120mm]{figure4.eps} \\caption{Factional rms amplitude as a function of the QPO centroid frequency for XTE J1701--462. Black points represent QPOs in the atoll phase, and grey points represent QPOs in the Z phase. Empty and filled squares indicate lower and upper kHz QPOs, respectively. Arrows correspond to 95\\% confidence level upper limits.} \\label{r_nu} \\end{center} \\end{figure*} The upper and the lower panels show respectively the maximum quality factor and the maximum rms amplitude of the lower kHz QPO for 13 sources as a function of the luminosity of the source in the 2-50 keV range, normalised by the Eddington luminosity, L$_{Edd}=2.5\\,\\times10^{38}$ erg s$^{-1}$, corresponding to a neutron star of 1.9 M$_{\\odot}$ accreting gas with cosmic abundance. As it is apparent from the plots, XTE J1701--462 follows the trends already traced by the other sources; moreover the coherence of the kHz QPOs in the Z phase fill the gap between the atoll sources in the left-hand side and the Z sources in the right-hand side of the graph, strengthening the correlation. (Notice, however that the luminosity of XTE J1701--462 could be uncertain by up to a factor of $\\sim$ 2, \\citealt{l2009a}).\\\\ Besides differences of the kHz QPOs properties between Z and atoll sources, Figure \\ref{max} shows also a trend of both $Q$ and rms amplitude within the atoll sources. This is noticeable in the lower panel of Figure 5 of \\citet{m2006}, where the maximum coherence and maximum rms amplitude of the lower kHz QPO are plotted vs. each other. From that plot it is apparent that the two quantities both in Z and atoll sources follow the same correlation \\citep[see][]{m2006}. This suggests that there is a single mechanism behind this trend. Our observations of XTE J170--462 in the Z and atoll phase are in line with this.\\\\ We can qualitatively explain how it is possible to find high rms amplitudes of kHz QPOs at energies where the contribution of the disk is insignificant. As reported by \\citet{b1996}, the rms amplitude of the lower kHz QPO in the LMXB 4U 1608--52 increases with energy up to 20\\% (fraction of the total emitted flux) at energies around 30 keV, significantly above the energy range where the accretion disk contributes to the X-ray spectrum (see also \\citealt*{mvf2001}; \\citealt{g2003}; \\citealt{a2008}). If the kHz QPOs are produced in the accretion disk, this must imply the presence of a mechanism that amplifies the variability at different energy bands. Some mechanisms have been proposed for this:  \\citet{lm1998} found that oscillations of the density and temperature of the Comptonizing medium can reproduce the rms amplitude behaviour of the lower kHz QPO in the atoll source 4U 1608--52. \\citet*{g2003}, from the analysis of the Z source GX340+0, suggest that QPOs are related to the contribution of the boundary layer emission to the total source emission. \\citet{gr2005} found that as $\\dot{m}$ increases from the horizontal branch to the flaring branch along the Z-shaped track in the colour-colour diagram of GX 340+0, the contribution of the boundary layer decreases. If we combine the results of \\citet*{g2003} and \\citet{gr2005}, and we apply them to the case of  XTE J1701--462, we would expect the boundary-layer contribution to be stronger in the atoll phase than the in Z phase. To verify this, we used the spectral analysis of XTE J1701--462 done by \\citet{l2009a}; their Figures 14 and 15 show the spectral fitting results, respectively, of the atoll and the Z phase. From their Figure 14 we notice that in the atoll phase the blackbody (BB) component used by \\citet{l2009a} to fit the boundary-layer emission becomes dominant, while the emission from the disk (fitted with a multi colour disk blackbody, MCD) becomes negligible. Comparing this with their Figure 15, we found that the fractional contribution of the boundary-layer emission in the atoll phase is higher than in the Z phase, which is in agreement with the results of \\citet*{g2003}, and can also explain the fact that we found stronger kHz QPOs in the atoll phase than in the Z phase. If this is correct, we should expect that most of the variability concentrates in the energy band where the boundary-layer emission peaks. From Figure 14 of \\citet{l2009a}, the BB temperature is about 2 keV, which implies that the peak of the boundary-layer emission is at about 10 keV. According to the results shown in the previous section, the atoll fractional rms amplitude increases as the energy increases, up to about 20\\% in the energy band 16-25 keV. We test whether the fractional rms amplitude we found above 10 keV is compatible with the fractional contribution of the boundary layer emission to the total emission of the source. From the spectral fits shown by \\citet[][see their Figure 12 panel ``atoll SS'']{l2009a}, the BB contributes more than 30\\% of the total emission in the 10-25 keV band, which means that the picture where the amplitude and coherence of the kHz QPOs are driven by the boundary layer is still consistent at those energies. Further studies should address the fact the kHz QPOs are only detected over narrow regions in the colour-colour diagram.\\\\ \\begin{figure} \\begin{center} {\\includegraphics[width=60mm,angle=-90]{figure5.eps}} \\caption{Upper panel: the maximum fractional rms amplitude of the lower kHz QPO for a sample of 12 sources (filled circles, see Table 1 \\citealt{m2006} for the source names) plus XTE J1701--462 as a function of the luminosity of the source. The diamond  and the triangle represent, respectively, measurements in the atoll and the Z phase. Lower panel: the maximum quality factor of the lower kHz QPO for the same sources mentioned above as a function of the luminosity of the source. The symbols are the same as in the upper panel. The luminosity is in units of the Eddington luminosity for a 1.9 M$_{\\odot}$ neutron star.} \\label{max} \\end{center} \\end{figure} Adding all these together, we suggest a possible scenario to describe the behaviour of the properties of the kHz QPO. Mathematically a Lorentzian (which is usually used to model quasi-periodic oscillations) is the Fourier transform of an exponentially damped sinusoid signal, $A\\,e^{-t/\\tau} sin (2\\pi \\nu t+\\phi)$, where $A$ is the amplitude of the signal, $\\tau$ is the life time of the variability, $\\nu$ is the frequency of the oscillation, and $\\phi$ is an arbitrary phase. Starting from this mathematical expression we can build a qualitative mechanism to describe the behaviour of the quality factor and the rms of the kHz QPO. Most of the models proposed to explain the kHz QPOs, consider the disk as the most probable location where those QPOs are generated (e.g. \\citealt*{mlp1998}). Under this assumption, the above expression could represent the oscillator that generates the QPO in the disk and sets its frequency $\\nu$. Since the rms amplitude of kHz QPOs depends on energy, we should introduce an energy-dependent amplification factor $B(E)$ which should be physically related to the mechanism which generates high-energy photons. For example, this could be related to the properties of the Comptonizing medium \\citep{lm1998} or to the contribution of the boundary-layer emission to the total emission of the source \\citep*{g2003, gr2005}. The life time $\\tau$ drives the QPO width, and as the amplitude, also this could be energy dependent. Additionally, we should consider that each process that amplifies the variability could also modify the life time of the variability. Just to give an example, if the QPO is created in the disk, and later on the QPO photons interact with matter in the corona, the final width of the QPO will be the result from the combination of the two process. In our model this would imply a lifetime for the QPO $1/\\tau_{tot}=(\\tau_{disk}+\\tau_{cor})/(\\tau_{disk}*\\tau_{cor})$. With these considerations, the modulated flux in the kHz QPO could be described by the expression $B(E)*A\\,e^{-t/\\tau_{tot}}\\,sin(2\\pi \\nu t+\\phi)$. This simple expression describes the rms amplitude and coherence of the QPO in terms of the quantities $B(E)$ and $\\tau_{tot}$, which in turn could depend on the properties of the boundary layer. The above expression would reproduce the behaviour of the rms amplitude and coherence if, for instance, $B(E)$ and $\\tau_{tot}$ depended upon mass accretion rate, $\\dot m$, such that $B(E)$ and $\\tau_{tot}$ decreased as $\\dot m$ increased. The previous description does not resolve however the issue of which processes are involved, or what the key mechanisms are that create or amplify the QPO signal. Nevertheless, this qualitative explanation provides a starting point to build a realistic model to explain the kHz QPO properties.\\\\ For most of the considerations in this paper we assumed that the only kHz QPO in the atoll phase is the lower kHz QPO. Although it is very unlikely (see section \\ref{res}), if it instead was the upper kHz QPO, one would have to explain the fact that in the frequency range, from 740 Hz to 860 Hz, the atoll kHz QPOs show $Q$ and rms values significantly different than those of the Z phase kHz QPOs (Figures \\ref{Q_nu} and \\ref{r_nu}).\\\\ From the INTEGRAL catalog (INTEGRAL general reference catalog, Version 30) we found no other sources within 1 square degree of the position of XTE J1701--462. The closest source reported (about 2.3 degrees away from XTE J1701--462) is the Gamma-ray source 2EGS J1653-4604 that does not show X-ray emission. \\citet{kr2006} observed XTE J1701--462 in outburst with the Chandra's High-Resolution Camera (HRC-S) in timing mode and they did not find evidence of other sources in the HRC-S field of view ($6' \\times 90'$). Swift observations with the X-ray telescope as well as XMM Newton observations of XTE J1701--462 did not show other sources, respectively within $23.6' \\times 23.6'$ and $30' \\times 30'$. Finally we checked the RXTE Galactic Bulge scans archive and we found that with a spatial accuracy of $15'$ no sources have been detected in a 1 square degree area from the XTE 1701-462 position. The evidence just discussed indicates that most likely the outburst observed by RXTE was due to a single source, XTE J1701--462, switching from Z to atoll.\\\\ Another interesting aspect to mention is the mismatch of the frequency of the lower kHz QPOs between the two different phases of the source. It is clear from Tables \\ref{tab2} and \\ref{tab1} and Figures \\ref{Q_nu} and \\ref{r_nu} that Z phase lower kHz QPOs are in the range $500-660$ Hz, while in the atoll phase the lower kHz QPOs are in the frequency range $640-860$ Hz. Although a similar effect has been observed when one compares other atoll and Z sources (see \\citealt{m2006} for a compilation) , we have now shown this effect in a single source. Studying the energy spectra in different parts of the colour-colour diagram, \\citet{l2009a} found that in the Z phase the disk is truncated far from the NS surface and its inner radius could be set by the local Eddington limit, while in the atoll phase the disk extends closer to the NS. Different sizes of the inner radius of the accretion disk could explain the different range of kHz QPO frequencies. However we can not discard the presence of other lower kHz QPOs in the Z phase at frequencies higher than what we found. From the upper limits for the Z phase reported in section \\ref{res}, it is apparent that we are not sensitive enough to detect broad and weak QPOs at high frequencies in that phase."
    },
    "1005/1005.2950_arXiv.txt": {
        "abstract": "{} {We report new Very Long Baseline Array (VLBA) polarimetric observations of the Compact Steep-Spectrum (CSS) sources 3C\\,119, 3C\\,318, and 3C\\,343 at 5 and 8.4~GHz. } {By using multifrequency VLBA observations we have derived milliarcsecond-resolution images of the total intensity, polarisation, and rotation measure ($RM$) distributions. } { The CSS source 3C\\,119, associated with a possible quasar, has source rest-frame $RM$ values up to $\\sim$10200 rad m$^{-2}$  in a region which coincides with a change in the direction of the inner jet.  This component is located $\\sim$325 pc from the core, which is variable and has a peaked radio spectrum.  In the case of 3C\\,318, associated with a galaxy, a rest-frame $RM$ of $\\sim$3030 rad m$^{-2}$ has been estimated for the brightest component which contributes almost all of the polarised emission. Further, two more extended components have been detected, clearly showing ``wiggles'' in the jet towards the southern side of the source. The CSS source 3C\\,343 contains two peaks of emission and a curved jet embedded in more diffuse emission. It exhibits complex field directions near the emission peaks, which indicate rest-frame $RM$ values in excess of $\\approx$6000 rad m$^{-2}$.  The locations of the cores in 3C\\,318 and 3C\\,343 are not clear. } {The available data on mas-scale rest-frame $RM$ estimates for CSS sources show that these have a wide range of values extending up to $\\sim$40000 rad m$^{-2}$ in the central region of OQ172, and could be located at projected distances from the core of up to $\\sim$1600 pc, as in 3C\\,43 where this feature has a rest-frame $RM$ of $\\sim$14000 rad m$^{-2}$.  $RM$ estimates for cores in core-dominated radio sources indicate that in addition to responding to an overall density gradient of the magneto-ionic medium, geometry, orientation and modes of fuelling may also play a significant role. In addition to these effects, the high values of $RM$ in CSS sources are possibly due to dense clouds of gas interacting with the radio jets. The observed distortions in the radio structures of many CSS sources are consistent with this interpretation.  } ",
        "introduction": "The number of Compact Steep-Spectrum (CSS) sources with detailed polarimetric information available at milliarcsecond resolution is still small. Polarised radio emission from CSS radio galaxies is either very weak or below the detection limits at centimetre wavelengths. In contrast, CSS quasars show linear polarisation percentages of up to 10\\% above 1\\,GHz \\citep [ ] [and references therein] {Rossetti08}.  We have conducted a series of observations of CSS sources having significantly polarised emission and high values of rotation measure ({\\it RM}) using the Very Long Baseline Array (VLBA).  CSS objects are {\\it young} radio sources, with ages $<10^{{\\rm 3-5}}$\\,yr. They have linear sizes $\\leq 20$\\,kpc \\footnote {$H_0=71\\,{\\rm km}\\, {\\rm s}^{-1}\\, {\\rm Mpc}^{-1}, \\Omega_{\\rm m}=0.27, \\Omega_{\\rm vac}=0.73$} and steep high-frequency radio spectra ($\\alpha >0.5$; ${\\rm S}_{\\nu}\\propto\\nu^{-\\alpha}$). Being sub-galactic in size, CSS sources reside deep within their host galaxies.  Therefore, Faraday rotation effects are to be expected when their polarised synchrotron emission is observed through the host galaxy magneto-ionic interstellar medium (ISM). The comparison of polarised emission over a range of wavelengths is an important diagnostic of the physical conditions within and around these compact radio sources (see Cotton et al. 2003c for an overview).  Existing sub-arcsec polarimetry has provided evidence in favour of the interaction of components of CSSs with dense clouds of gas, as seen for example in the CSS quasar 3C\\,147 \\citep {Junor99a}.  Results for the first two CSS quasars observed in our on-going program, \\object{B$0548+165$} and \\object{B$1524-136$}, are available in  \\citet{Mantovani02}, while those for \\object{3C\\,43} (B$0127+233$) are to be found in \\citet{Mantovani03}. More recently, the results for \\object{3C\\,147} (\\object{B$0538+498$}) have been presented by ~\\citet{Rossetti09}.  These sources have all been imaged with milliarcsecond resolution by means of full-Stokes VLBA observations.  In this paper, we report on multi-frequency VLBA, plus a single Very Large Array (VLA) antenna, polarisation observations at 5 and 8.4\\,GHz for \\object{3C\\,119} (\\object{B$0429+415$}), \\object{3C\\,318} (\\object{B$1517+204$}), and \\object{3C\\,343} (\\object{B$1634+628$}).  In Section \\ref{sec:observation} we summarise the observations and data processing. Section \\ref{sec:sources} describes the new information obtained on the structural and polarisation properties of \\object{3C\\,119, 3C\\,318, and 3C\\,343}. Discussion and conclusions are presented in Sections \\ref{sec:discussion}  and \\ref{sec:conclusions} respectively. ",
        "conclusions": "\\label{sec:conclusions} We have presented multi-frequency VLBA polarisation observations of three CSS sources, namely 3C\\,119, 3C\\,318 and 3C\\,343 to estimate their {\\it RM} values. The radio source 3C\\,119 is associated with a possible quasar, and its $RM$ in the source rest frame has been found to be as high as $\\sim$10200 rad m$^{-2}$ in a region which coincides with a change in direction of the inner jet.  This component is located at a projected distance of $\\sim$325 pc from the core, which is almost point-like, variable, has a peaked radio spectrum and is at best marginally polarised. 3C\\,318 is associated with a radio galaxy and its rest frame {\\it RM} has been found to reach a maximum of $\\sim$3030 rad m$^{-2}$ for the brightest component which contributes almost all of the polarised emission. These observations are more sensitive than those of Spencer et al. (1991) and have detected two more extended components, which trace  ``wiggles'' in the jet towards the southern side of the source.  Of the three, the CSS source 3C\\,343 has perhaps the most complex structure. It contains two peaks of emission and a curved jet embedded in more diffuse emission. It exhibits complex field directions near the peaks of emission, which indicate rest frame $RM$ values in excess of $\\approx$4000 rad m$^{-2}$.  The varying sensitivity for the different frequencies and the complex field patterns near the peaks of emission make it difficult to construct a reliable {\\it RM} image for this source.  We have compiled the available data on mas-scale $RM$ estimates for CSS sources. These show a wide range of values with indications of a low {\\it RM} for 3C\\,287 which needs to be confirmed from observations with a larger number of frequencies, to  values as high as $\\approx$40000 rad m$^{-2}$ in the central region of OQ172 (Udomprasert et al. 1997).  The components with high {\\it RM} can also occur at considerable distances from the core, e.g. in 3C\\,43 where the component with an {\\it RM} of $\\sim$14000 rad m$^{-2}$ is located at a projected distance of $\\sim$1600 pc (Cotton et al. 2003a).  {\\it RM} estimates for flat-spectrum cores in largely core-dominated radio sources appear to increase with frequency (see O'Sullivan \\& Gabuzda 2009), suggesting that as one probes deeper into the core or unresolved base of the jet, one samples regions of higher density and/or magnetic field in the magneto-ionic medium. On larger scales, the jet {\\it RM}s tend to be low as these objects are observed along a line of sight where the magneto-ionic medium may have been swept out by the relativistic jets (e.g. Saikia et al. 1998; Taylor 2000; Zavala \\& Taylor 2004).  The CSS objects for which {\\it RM} values have been estimated are almost entirely quasars.  While the effects of an overall density gradient in the magneto-ionic medium, along with effects of geometry, orientation {\\bf and modes of fuelling of the AGN} , are likely to play a significant role, the high $RM$ values in many of these CSS sources appear to be due to dense clouds of gas interacting with the radio jets.  Usually, they also exhibit large structural bends and distortions, consistent with the possibility of jet-cloud interactions in the interstellar medium of the host galaxy.  Some of these gas clouds may also be responsible for fuelling the AGN activity."
    },
    "1005/1005.5447.txt": {
        "abstract": "A new fully quantum method describing penetration of packet from internal well outside with its tunneling through the barrier of arbitrary shape used in problems of quantum cosmology, is presented. The method allows to determine amplitudes of wave function, penetrability $T_{\\rm bar}$ and reflection $R_{\\rm bar}$ relatively the barrier (accuracy of the method: $|T_{\\rm bar}+R_{\\rm bar}-1| < 1 \\cdot 10^{-15}$), coefficient of penetration (i.~e. probability of the packet to penetrate from the internal well outside with its tunneling), coefficient of oscillations (describing oscillating behavior of the packet inside the internal well). % Using the method, evolution of universe in the closed Friedmann--Robertson--Walker model with quantization in presence of positive cosmological constant, radiation and component of generalize Chaplygin gas is studied. It is established (for the first time): (1) oscillating dependence of the penetrability on localization of start of the packet; (2) presence of resonant values of energy of radiation $E_{\\rm rad}$, at which the coefficient of penetration increases strongly. % From analysis of these results it follows: (1) necessity to introduce initial condition into both non-stationary, and stationary quantum models; (2) presence of some definite values for the scale factor $a$, where start of expansion of universe is the most probable; (3) during expansion of universe in the initial stage its radius is changed not continuously, but passes consequently through definite discrete values and tends to continuous spectrum in latter time. % % (4) \u00f1\u00eb\u00e5\u00e4\u00f1\u00f2\u00e2\u00e8\u00e5\u00ec \u00ea\u00e2\u00e0\u00ed\u00f2\u00ee\u00e2\u00e0\u00ed\u00e8\u00ff \u00ea\u00ee\u00f1\u00ec\u00ee\u00eb\u00ee\u00e3\u00e8\u00f7\u00e5\u00f1\u00ea\u00ee\u00e9 \u00ec\u00ee\u00e4\u00e5\u00eb\u00e8 \u00ff\u00e2\u00eb\u00ff\u00e5\u00f2\u00f1\u00ff \u00e4\u00e8\u00f1\u00ea\u00f0\u00e5\u00f2\u00ed\u00ee\u00f1\u00f2\u00fc \u00ef\u00f0\u00ee\u00f1\u00f2\u00f0\u00e0\u00ed\u00f1\u00f2\u00e2\u00e0-\u00e2\u00f0\u00e5\u00ec\u00e5\u00ed\u00e8, \u00ef\u00f0\u00ee\u00ff\u00e2\u00eb\u00ff\u00fe\u00f9\u00e0\u00ff\u00f1\u00ff \u00ed\u00e0\u00e8\u00e1\u00ee\u00eb\u00e5\u00e5 % \u00f1\u00e8\u00eb\u00fc\u00ed\u00ee \u00e2 \u00ef\u00e5\u00f0\u00e2\u00ee\u00e9 \u00f1\u00f2\u00e0\u00e4\u00e8\u00e8 \u00f0\u00e0\u00f1\u00f8\u00e8\u00f0\u00e5\u00ed\u00e8\u00ff \u00e2\u00f1\u00e5\u00eb\u00e5\u00ed\u00ed\u00ee\u00e9. ",
        "introduction": "}  If we analyzed existed variety of quantum models which describe formation of the universe and its subsequent evolution in the first stage, we should come to conclusion that the semiclassical approach for description of tunneling and determination of wave function is the most prevailing today. This approach forms a basis, which props up both models with the Feynman formalism of path integrals in multidimensional space-time, developed by the Cambridge group and other researchers, called the \\emph{``Hartle--Hawking method''} (for example, see Ref.~\\cite{Hartle.1983.PRD}), and methods based on direct consideration of tunneling in 4-dimensional Euclidian space-time called the \\emph{``Vilenkin method''} (for example, see Refs.~\\cite{Vilenkin.1982.PLB,Vilenkin.1983.PRD,Vilenkin.1984.PRD,Vilenkin.1986.PRD,Vilenkin.1988.PRD,% Vilenkin.1994.PRD,Vilenkin.1995,Vilenkin.1997.PRD.Comments,Vilenkin.2003.PRD}). The models in 4-dimensional space-time directed on description of inflation, on study of fluctuations of vacuum with inclusion of massive fields (for example, see~Refs.\\cite{Bouhmadi-Lopez.2002.PRD}), multidimensional cosmological models (for example, see~Refs.\\cite{Kolb.1990,Carugno.1996.PRD}), variety of string models (for example, see~Refs.\\cite{Brustein.2006.PRD,Kobakhidze.2007.EPJC}) have mainly such a semiclassical grounds. To date, this basis is supposed to be sufficiently reliable and give interested quantum characteristics of the universe with good accuracy. Such a point of view so has been taken root and is prevailing, that, in spite of almost 40 years of researches in quantum cosmology (see~the first papers~\\cite{DeWitt.1967,Wheeler.1968}), the papers devoted on more tiny and deeper study of possible quantum nature of the universe, its formation and evolution on the first stage, can be found enough rarely (for example, see \\cite{Rubakov.2002.PRD,AcacioDeBarros.2007.PRD,Monerat.2007.PRD}, also~\\cite{Yurov.2004.TMP,Garcia.2006.IJTP}).  However, one should ask whether the penetrability determined according to the semiclassical theory by a shape of the barrier solely between two turning points, gives exhaustive answer and the best estimations of rates of evolution of universe.  % \\vspace{2mm} (1) If, despite such widespread confidence in the semiclassical approach, we still want to check it, we immediately will miss some of the parameters. For example, it would seem, one can use a test of $T + R = 1$ (where $T$ and $R$ are penetrability and reflection relatively the barrier in the cosmological problem), where usually we have no doubt. However, one should recall that in quantum mechanics, the semiclassical approximation neglects the reflected waves completely (see~\\cite{Landau.v3.1989}, eq.~(46.10), p.~205, also p.~221--222) and, therefore, to compare the calculated penetrability $T$ is just not with anything.  % \\vspace{2mm} (2) If we still wanted to determine the reflection coefficient, then we should increase the order of approximation of the semiclassical method (in order to take into account the decreasing component of the wave function on the background of increasing one in the region of tunneling), and just stumble on the next problem --- presence of a non-zero interference between the incident and reflected waves in the radial task. Now the criterion $T + R = 1$ for testing is not satisfied and it needs to take into account the third component $M$ of interference in addition (see~\\cite{Maydanyuk.2010.IJMPD}). In particular, at unsuccessfully chosen separation of the exactly known full wave function on the incident and reflected waves (but the semiclassical approaches have no needed apparatus for such analysis), the interference component can increase without limit, and be substantially larger than the penetrability and reflection. After the appearance of such an arbitrariness, the penetrability and reflection can freely exceed the unit and increase without limit. In what is now the general meaning of the penetrability?  % \\vspace{2mm} (3) We shall give only some easy examples from quantum mechanics. (i) If we consider two-dimensional penetration of the packet through the simplest rectangle barrier (with finite size), we shall see that the penetrability is directly dependent on direction of tunneling of the packet. So, the penetrability is not a single value but the function. (ii) If to consider one-dimensional tunneling of the packet through the simplest rectangular barrier, one can obtain ``interference picture'' of its amplitude in the transmitted region, which is depended on time and space coordinates and is an exact analytical solution. Of course, the stationary part of such a result coincides exactly with well known stationary solutions \\cite{Maydanyuk.2003.PhD-thesis}. From the arguments above the impression could appear that the penetrability defined solely by the shape of the barrier between two turning points is nothing more than prevailing simplified understanding, while for more accurate and deep analysis we need in the strong basis.  % \\vspace{2mm} (4) Advance of the semiclassical approach is in simplicity of formula of the penetrability based on determination of the outgoing wave in the asymptotic region. A \\emph{tunneling boundary condition} \\cite{Vilenkin.1995,Vilenkin.1994.PRD} seems to be natural and clear, where the wave function should represent an outgoing wave at large scale factor $a$. However, whether is such a wave free? In contrast to problems of quantum atomic and nuclear physics, in cosmology we deals with potentials, which modules only increase with increasing the scale factor $a$ (also their gradients increase, which have sense of force acting on the wave) and, therefore, we have nothing mutual with a free propagation of the wave in the asymptotic region. Now it is unclear to which combination two Airy functions should be combined at turning point, in order to obtain the proper outgoing wave. It turns out that instead of the free wave in the asymptotic (missing in problems of quantum cosmology), we should be able to work with the waves propagating inside strong fields (see~\\cite{Maydanyuk.2010.IJMPD}).  % \\vspace{2mm} These problems violate (destroy) the basis of the semiclassical models, and now statements about reliability of the semiclassical models are transformed into the question of `` faith'' in them, though widespread \\cite{Maydanyuk.2010.IJMPD}. The semiclassical approach could be compared to \\emph{``black box''} in which deeper and more detailed information about the dynamics of the universe is hide. More importantly, in such a black box those missing elements are hidden, without which it is impossible to combine everything together to obtain self-consistent formalism of quantum description of the formation of the universe and its evolution in the first stage. In order to clarify these questions, we have developed a new fully quantum method presented in this paper.  % \\vspace{2mm} This paper is organized so. In Sec.~2 a new fully quantum method for description of the formation of the universe and its evolution in the first stage on the basis of a packet, which penetrates from the internal potential well outside by tunneling through the barrier, is presented. A main advance of this method is determination of characteristics of the packet with high accuracy (without implication of the semiclassical approximations). The formalism of the method has been developed relatively the barrier of arbitrary shape, that makes the method universal. In Sec.~3 the method is applied for solution of the problem of evolution of the packet in FRW-model with radiation and Chaplygin gas. The penetrability and reflection relatively the barrier are calculated. We propose new characteristics, more adequately determining the probability of formation of the universe and its subsequent expansion. An accuracy of the method is demonstrated, achieving to $|T_{\\rm bar}+R_{\\rm bar}-1| < 1 \\cdot 10^{-15}$ inside whole under-barrier range of the energy of radiation ($M = 0$, author has not yet found competitive approaches, achieving such a precision), stability in calculations for the obtained results is shown. A special attention in analysis is devoted to study of the initial conditions. On such a basis, for the first time oscillatory dependence of the penetrability on the localization of start of the packet and resonant levels of the energy of radiation $E_(\\rm rad)$ (where the penetration extremely increases) are opened (which are hidden in the semiclassical picture). In finishing, in Sec.~\\ref{sec.conclusions} conclusions of the obtained results are formulated. % *******************************************************************************************************************    % ******************************************************************************************************************* ",
        "conclusions": "}  The new method for determination of probability of penetration of the packet from the internal well outside with its tunneling through one-dimensional barrier of arbitrary shape used in problems of quantum cosmology, is presented. Note the following: % \\begin{enumerate} \\item The method is further development of approach of multiple internal reflections (see Refs.~\\cite{Maydanyuk.2000.UPJ,Maydanyuk.2002.JPS,Maydanyuk.2002.PAST,Maydanyuk.2003.PhD-thesis,% Maydanyuk.2006.FPL,Maydanyuk.arXiv:0805.4165,Maydanyuk.arXiv:0906.4739}, also Refs.~\\cite{Fermor.1966.AJPIA,McVoy.1967.RMPHA,Anderson.1989.AJPIA,Esposito.2003.PRE}), where a process of tunneling of the packet through the barrier is considered consequently by steps of its propagation relatively to each boundary of the barrier. The method is fully quantum, allows to determine amplitudes of wave function, coefficients of penetrability and reflection relatively the barrier. For the first time \\emph{exact} analytical solutions for amplitudes of the wave function, penetrability $T$ and reflection $R$ for the barrier, composed from arbitrary number $n$ of rectangular potential steps, are found. At limit $n \\to \\infty$ these solutions could be considered as exact limits for potential with interested barrier and internal well of arbitrary shape.  \\item Accuracy of the method in determination of penetrability $T_{\\rm bar}$ and reflection $R_{\\rm bar}$ is: $|T_{\\rm bar}+R_{\\rm bar}-1| < 1 \\cdot 10^{-15}$ (see Tabl.~\\ref{table.model_Monerat.1}). Author has not found other methods achieving such accuracy in similar problems of quantum physics (with possible exception of some selected cases of exactly solvable barriers which could be obtained by methods of supersymmery).  \\item On the basis of the method the probability of penetration of the packet from the internal well outside with its tunneling through the barrier of arbitrary shape called \\emph{coefficient of penetration}, is determined. It succeeds to separate that coefficient explicitly on the penetrability and new coefficient, which characterizes oscillating behavior of the packet inside the internal well and is called \\emph{coefficient of oscillations}. That found for the first time formula seems to be fully quantum analogue of the semiclassical formula of $\\Gamma$ width of decay in quasistationary state proposed by Gurvitz and K\\\"{a}lbermann in Ref.~\\cite{Gurvitz.1987.PRL} (here, the coefficient of oscillations is fully quantum analogue for the semiclassical $F$ factor of formation and the coefficient of penetration is analogue for the semiclassical $\\Gamma$ width). \\end{enumerate}  This method has been applied for study of properties of the packet, describing evolution of universe on the first stage in the closed Friedmann--Robertson--Walker model with quantization in the presence of the positive cosmological constant, radiation and component of generalize Chaplygin gas with potential chosen from \\cite{Monerat.2007.PRD}. Let us formulate main results obtained for the first time: % \\begin{enumerate} \\item For the same chosen energy of radiation $E_{\\rm rad}$ the penetrability of the barrier is changed visibly in dependence on the position of the starting point $R_{\\rm start}$ inside the internal well, where the packet begins to propagate (see Fig.~\\ref{fig.model_Monerat.2}): the penetrability has oscillating behavior, difference between its minimums and maximums is minimal at $R_{\\rm start}$ in the center of the well, with increasing $R_{\\rm start}$ this difference increases achieving maximum close to the turning point. The behavior of the coefficients of reflection, oscillations and penetration turns out to be similar. Coincidence (up to the first 15 digits) of the amplitudes of the wave function obtained by such a method, with corresponding amplitudes obtained in the standard approach of quantum mechanics (see App.~\\ref{sec.app.1}) at different energies $E_{\\rm rad}$ confirms that this result does not depend on a choice of the fully quantum method applied for calculations. Such a peculiarity is shown in the fully quantum non-stationary and stationary considerations and it is hidden after imposing the semiclassical restrictions.  \\item In non-stationary and stationary considerations the penetrability of the barrier should be connected with initial condition localizing start of the packet. Note that possible introduction of the initial condition into known stationary semiclassical models could change their results.  \\item The penetrability is changed visibly, if to take the external tail of the barrier into account. For example, for the barrier (\\ref{eq.model.2.2}) with parameters $A=0.001$ and $B=0.001$ (see Fig.~\\ref{fig.model_Monerat.1}) at the energy of radiation $E_{\\rm rad} = 223$ the penetrability is changed up to 2 times (see Fig.~\\ref{fig.model_Monerat.3}). If to increase the external boundary $a_{\\rm max}$, all amplitudes and coefficients are convergent in calculations that confirms efficiency of the developed method.  \\item Dependence of the coefficient of penetration on the energy of radiation has oscillating behavior: here peaks are clearly shown, localized approximately on the same distances, between which smooth minimums (wells) stable in calculations are observed (see Fig.~\\ref{fig.model_Monerat.4}). By other words, for the first time in the fully quantum approach we have obtained clear and stable picture of resonances, which indicate on presence of some quasistationary states. Here, with increasing of the energy of radiation the penetrability is changed monotonously and it describes a general tendency of behavior of the coefficient of penetration, while the coefficient of oscillations gives peaks. Now a reason of existence of resonances becomes clear --- oscillations of the packet inside the internal well give them. For example, for the barrier (\\ref{eq.model.2.2}) with parameters $A=0.001$ and $B=0.001$ I establish 134 such resonant levels inside range $E_{\\rm rad}$ = 0--223 (see Tabl.~\\ref{table.model_Monerat.2},~\\ref{table.model_Monerat.3}).  \\item The dependence of the penetrability on the starting point has clear established maximums and minimums. On such a basis one can suppose that the most probable start of the packet (describing start of expansion of the universe) is in one point of such maximums. This allows to predict some definite initial values of the scale factor, at which the universe begins to expand (such initial data is direct result of quantization of the classical cosmological model).  \\item Modulus of the wave function both in the internal, and in the external regions has clear established own maximums and minimums~\\cite{Maydanyuk.2008.EPJC,Maydanyuk.2010.IJMPD}. This indicates on such values of the scale factor, at which ``appearance'' of the universe will be more or less probable. By other words, radius of the universe during its expansion is changed not continuously, but passes consequently through some definite discrete values connected with these maximums. It follows that after quantization space-time of universe on the first stage of its expansion seems to be rather discrete than continuous. According to results~\\cite{Maydanyuk.2008.EPJC,Maydanyuk.2010.IJMPD}, difference between maximums and minimums with increasing of the scale factor $a$ is slowly smoothed and we obtain obvious for us continuous structure of the space-time at latter times. Discontinuity of space-time is direct result of quantization of cosmological model, which is shown the most strongly on the first stage of expansion and disappeared after imposition of the semiclassical approximations. \\end{enumerate} % *******************************************************************************************************************     % ******************************************************************************************************************* \\appendix"
    },
    "1005/1005.0036_arXiv.txt": {
        "abstract": "\\hspace{\\parindent} {\\HII} regions are the birth places of stars, and as such they provide the best measure of current star formation rates (SFRs) in galaxies.  The close proximity of the Magellanic Clouds allows us to probe the nature of these star forming regions at small spatial scales.  To study the {\\HII} regions, we compute the bolometric infrared flux, or total infrared (TIR), by integrating the flux from 8 to 500~$\\mu$m.  The TIR provides a measure of the obscured star formation because the UV photons from hot young stars are absorbed by dust and re-emitted across the mid-to-far-infrared (IR) spectrum.  We aim to determine the \\textit{monochromatic} IR band that most accurately traces the TIR and produces an accurate obscured SFR over large spatial scales.  We present the spatial analysis, via aperture/annulus photometry, of 16 Large Magellanic Cloud (LMC) and 16 Small Magellanic Cloud (SMC) {\\HII} region complexes using the \\textit{Spitzer Space Telescope's} IRAC (3.6, 4.5, 8~$\\mu$m) and MIPS (24, 70, 160~$\\mu$m) bands.  Ultraviolet rocket data (1500 and 1900 \\AA) and SHASSA {\\Ha} data are also included.  All data are convolved to the MIPS 160~$\\mu$m resolution (40 arcsec full width at half-maximum), and apertures have a minimum radius of $35\\arcsec$. The IRAC, MIPS, UV, and {\\Ha} spatial analysis are compared with the spatial analysis of the TIR.  We find that nearly all of the LMC and SMC {\\HII} region spectral energy distributions (SEDs) peak around 70~$\\mu$m at all radii, from $\\sim10$ to $\\sim400$ pc from the central ionizing sources.  As a result, we find the following: the sizes of {\\HII} regions as probed by 70~$\\mu$m is approximately equal to the sizes as probed by TIR ($\\approx70$ pc in radius); the radial profile of the 70~$\\mu$m flux, normalized by TIR, is constant at all radii (70~$\\mu\\mbox{m}\\sim0.45~\\mbox{TIR}$); the $1\\sigma$ standard deviation of the 70~$\\mu$m fluxes, normalized by TIR, is a lower fraction of the mean ($0.05-0.12$ out to $\\sim220$ pc) than the normalized 8, 24, and 160~$\\mu$m normalized fluxes ($0.12-0.52$); and these results are the same for the LMC and the SMC.  From these results, we argue that 70~$\\mu$m is the most suitable IR band to use as a monochromatic obscured star formation indicator because it most accurately reproduces the TIR of {\\HII} regions in the LMC and SMC and over large spatial scales.  We also explore the general trends of the 8, 24, 70, and 160~$\\mu$m bands in the LMC and SMC {\\HII} region SEDs, radial surface brightness profiles, sizes, and normalized (by TIR) radial flux profiles.  We derive an obscured SFR equation that is modified from the literature to use 70~$\\mu$m luminosity, $\\mbox{SFR}~(M_{\\odot}~\\mbox{yr}^{-1}) = 9.7(0.7)\\times10^{-44}~L_{70}~(\\mbox{ergs}~\\mbox{s}^{-1})$, which is applicable from 10 to 300 pc distance from the center of an {\\HII} region.  We include an analysis of the spatial variations around {\\HII} regions between the obscured star formation indicators given by the IR and the unobscured star formation indicators given by UV and {\\Ha}.  We compute obscured and unobscured SFRs using equations from the literature and examine the spatial variations of the SFRs around {\\HII} regions. ",
        "introduction": "\\hspace{\\parindent} {\\HII} regions are locations of active or recent star formation where the extreme ultraviolet (UV) radiation from massive OB stars ionizes the surrounding gas \\citep[for a good review see chapter 5 of][and references therein]{tiel05}.  It is common to consider {\\HII} regions as being large complexes of overlapping individual {\\HII} regions where the sizes and components are not determined solely by the classical Str\\\"{o}mgren radius of an individual star but by an inclusion of all of the physical regimes in the interstellar medium (ISM) affected by the far-UV photons of the central sources.  These {\\HII} region components include the central ionizing OB stars, and the surrounding photodissociation regions (PDRs) and molecular clouds, where the chemistry and the heating are driven by the UV photons from nearby stars \\citep{tiel05,rela09,wats08}.  Thus, {\\HII} regions in this context are really {\\HII} complexes, or more generally, star forming regions.  The observed properties of these star forming regions, including size and shape, depend on the physical traits such as the number of young hot ionizing stars, density of neutral gas, abundance of dust, star formation history, and prior supernovae \\citep[e.g., see][]{hodg74,tiel05,walb02,snid09,harr09}.  These physical traits are often correlated, and the use of data at many wavelengths to observe {\\HII} regions in different galactic environments is required to piece them together.  The nature of the hot central OB stars can be studied using UV \\citep[e.g.,][]{smit87,mart05} or nebular emission lines, such as {\\Ha} \\citep[e.g.,][]{heni56,hodg83,gaus01}.  The {\\Ha} luminosity function of a galaxy is frequently used to determine many of the physical properties of {\\HII} regions, including the number of ionizing stars, evolutionary effects, and possible environmental effects \\citep{kenn89,oey98}.  {\\HII} region size distributions are directly related to the luminosity functions and relate to the numbers of nebula of a given size for a galaxy\\citep[e.g.,][]{van81,oey03}.  The PDRs can be studied by many methods including the [\\OI] and [\\CII] infrared (IR) cooling lines and CO radio data \\citep[e.g.,][]{kauf99}. The molecular clouds are typically studied using radio observations of molecular rotational lines \\citep[e.g.,][]{cohe88,genz91,fuku08}.  Present star formation rates (SFRs) are calculated using a tracer of the UV photons from the young massive stars and spectral synthesis models \\citep[see][]{kenn98}.  {\\HII} regions will have some fraction of their UV photons obscured by dust and some fraction unobscured. For unobscured SFRs, the ionizing photons can be directly observed via UV observations or recombination lines such as {\\Ha}.  For obscured {\\HII} regions, bolometric IR observations of dust (i.e., the total infrared (TIR)) can be used to recover the extinguished UV photons. This is because the dust absorption cross section is highly peaked in the UV, and the re-emitted flux is in the broad spectral range from the mid-to-far-IR \\citep{kenn98}.  Because the TIR around {\\HII} regions accounts for all of the extincted UV photons, the TIR is expected to be the single best indicator of SFR obscured by dust.  Combining UV, optical, and IR observations of {\\HII} regions across a whole galaxy allows us to probe the current SFR of that galaxy.  The Large Magellanic Cloud (LMC) and Small Magellanic Cloud (SMC) are ideal natural laboratories for studying star forming regions and their effects on the ISM.  The close proximities of the LMC and SMC, at $\\sim52$ kpc \\citep{szew08} and $\\sim60$ kpc \\citep{hild05}, respectively, allow for detailed star formation studies down to parsec or sub-parsec scales depending on wavelength.  Any broad study of {\\HII} regions across the LMC and SMC can take advantage of many multiwavelength observations, including the rocket UV data from \\citet{smit87}, the Southern {\\Ha} Sky Survey Atlas (SHASSA) {\\Ha} data from \\citet{gaus01}, and the \\textit{Spitzer Space Telescope} IR data from \\citet{meix06} and {\\gord}.  Furthermore, there are many past optical {\\HII} region surveys that catalog the sources of {\\Ha} in the LMC and SMC \\citep[see][]{heni56,davi76,bica95}.  Because of the observational advantages, many researchers rely on a single-band star formation indicator (i.e., UV, {\\Ha}, {\\Pa}, 8~$\\mu$m, 24~$\\mu$m, etc).  There are complications in the UV and optical lines in that they can be greatly impacted by extinction, thus, requiring extinction corrections.  Another complication is that the observed UV photons can come from stars of various ages ($< 100$ Myr) and will greatly depend on galaxy type (i.e., quiescent spiral galaxies, starbursts, etc.; \\citep{calz05}).  The 8 and 24~$\\mu$m IR band emission will likely depend on the environment of the host galaxy because the abundance of the aromatics/small grains that give rise to their emission depend upon the metallicity and ionizing radiation present \\citep[K. D. Gordon et al. 2010, in preparation;][]{drai07,drai07b}.  Work done by \\citet{calz07} and \\citet{dale05} indicate that 8 $\\mu$m makes for a poor star formation indicator due to large variability of emission in galaxies with respect to spectral energy distribution (SED) shape and metallicity. \\citet{calz07} also note, along with \\citet{calz05}, that a star formation indicator using 24~$\\mu$m by itself can vary from galaxy to galaxy.  \\citet{dale05} claim that SFRs calculated from 24~$\\mu$m emission may be off by a factor of 5 due to variations of 24~$\\mu$m flux with respect to SEDs observed across nearby galaxies.  To compensate for the extinction effects in UV and optical nebular emission lines, many researchers are now measuring SFRs via a combination of obscured (TIR, 8~$\\mu$m, 24~$\\mu$m) and unobscured (UV, {\\Ha}, {\\Pa}) star formation indicators \\citep[e.g.,][]{calz07,kenn07,thil07,rela09,kenn09}.  However, the noted differences in 8 and 24~$\\mu$m emission, relative to host galaxy properties, may still introduce uncertainties to the calculated SFRs when applying them across large galaxy samples.  In their analysis of 33 galaxies from the \\textit{Spitzer} SINGS sample, \\citet{calz07} claim that a combination of {\\Ha} and 24~$\\mu$m gives the most robust SFR using a procedure similar to the \\citet{gord00} ``flux ratio method''\\footnote{The ``flux ratio method'' uses UV and IR fluxes of galaxies to derive extinction-corrected UV luminosities \\citep{gord00}}.  Their calibration has a caveat in that it is useful for actively star forming galaxies where the energy output is dominated by young stellar populations \\citep{calz07}.  \\citet{kenn09} analyze SFRs of nearby galaxies derived by combining {\\Ha} with 8~$\\mu$m, 24~$\\mu$m, and the TIR.  They find that linear combinations of {\\Ha} and TIR provide for the most robust SFRs.  There is little work done on investigating the efficacy of using 70 or 160~$\\mu$m as star formation indicators.  \\citet{dale05} claim that the 70~$\\mu$m emission may make a good monochromatic obscured star formation indicator because the 70 to 160~$\\mu$m ratio correlates well with local SFRs.  A \\textit{Spitzer} analysis of far-IR compact sources in the LMC \\citep{jacc10} and SMC \\citep{jacc10b}, including compact {\\HII} regions, find that the bolometric correction to 70 $\\mu$m is modest, due to a typical dust temperature of 40~K.  In a study of dwarf irregular galaxies, \\citet{walt07} find a good correlation between the brightest 70~$\\mu$m regions and optical tracers of star formation, albeit, with some galaxies contributing significant 70~$\\mu$m diffuse emission at large radii.  Which of the \\textit{Spitzer} IR bands most accurately reproduces the TIR over a large spatial scale?  Employing aperture/annulus photometry, we analyze 16 LMC and 16 SMC {\\HII} region complexes using the \\textit{Spitzer} Infrared Array Camera (IRAC) and Multiband Imaging Photometer (MIPS) bands.  We determine that the MIPS 70~$\\mu$m band provides for the most accurate monochromatic obscured star formation indicator based on an analysis of the {\\HII} region complex SEDs, sizes, and radial monochromatic IR fluxes (normalized by the TIR).  We include an analysis of the spatial distribution of the unobscured star formation indicators, UV and {\\Ha}, relative to the IR obscured star formation indicators.  We modify an established TIR SFR recipe from \\citet{kenn98} to derive a new monochromatic obscured SFR equation using the 70~$\\mu$m luminosity.  In Section~\\ref{sample}, we list the LMC and SMC {\\HII} region complexes sampled in this work.  In Section~\\ref{obs}, we explain the IR, UV, and {\\Ha} observations and data reduction.  The analysis of the multiwavelength photometry is discussed in Section~\\ref{analysis}. The basic results of the photometry, SEDs, and radial profiles, are discussed in Section~\\ref{results} as well as a discussion of our calculation of the TIR.  Discussions of {\\HII} region normalized radial SEDs, {\\HII} complex sizes, normalized radial profiles, and SFRs are presented in Section~\\ref{disc}.  In this section we also present our derived 70~$\\mu$m obscured SFR equation. We finish with some concluding statements in Section~\\ref{conc}. ",
        "conclusions": "\\label{conc} \\hspace{\\parindent}We performed UV, {\\Ha}, and IR aperture/annulus photometry of 16 LMC and 16 SMC {\\HII} regions at the common resolution of the MIPS 160~$\\mu$m $40\\arcsec$ {\\fwhm} PSF.  With the aperture photometry, we computed IR SEDs using the total cumulative \\textit{Spitzer} IRAC (3.6, 4.5, 8~$\\mu$m) and MIPS (24, 70, 160~$\\mu$m) fluxes of the LMC, SMC, and each {\\HII} region.  The cumulative SED of the LMC is brighter than the SMC, and their shapes are consistent with IR SEDs of the Magellanic Clouds in the literature.  The SMC {\\HII} regions have cumulative fluxes that are weaker, by about an order of magnitude, than the LMC {\\HII} regions. We argue that this is expected from what is known about the LMC and SMC {\\HII} region luminosity functions.  We compute the IR bolometric flux (TIR) for each {\\HII} region aperture/annulus by integrating under the linear extrapolation of the 8, 24, 70, and 160~$\\mu$m fluxes.  We include the input from the Rayleigh--Jeans tail of a computed dust MBB by integrating in steps of 0.25~\\AA, from 160 to $500~\\mu$m.  Using annulus photometry of the {\\HII} regions, at distances of $\\sim10$ up to $\\sim400$ pc from their cores, we plot the radial 8~$\\mu$m, 24~$\\mu$m, 70~$\\mu$m, 160~$\\mu$m and TIR surface brightness profiles of the 32 {\\HII} regions.  The surface brightness profiles peak near the center of the {\\HII} regions for all IR bands and TIR. We argue that resolution effects may play a role in reducing the SMC {\\HII} region surface brightness profiles relative to the LMC {\\HII} region surface brightness profiles.  The surface brightness profiles are used to compute the sizes of each {\\HII} region, using the TIR, 8~$\\mu$m, 24~$\\mu$m, 70~$\\mu$m, and 160~$\\mu$m emission.  The typical sizes measured have large scatter and are $\\approx75$ pc for 8~$\\mu$m, $\\approx50$ pc for 24~$\\mu$m, $\\approx70$ pc for both 70~$\\mu$m and TIR, and $\\approx90$ pc for 160~$\\mu$m.  The large scatter for each size measurement likely stems from sampling different parts of the LMC and SMC luminosity functions.  We compute normalized (by TIR) IR SEDs for each {\\HII} region annulus. The radial SEDs nearly all peak around 70~$\\mu$m at all radii for every {\\HII} region, out to $\\sim400$ pc.  The result of this is that it gives four favorable reasons to choose 70~$\\mu$m as a monochromatic obscured star formation indicator, and they are (1) 70~$\\mu$m emission most closely traces the size of {\\HII} regions as found using the TIR ($\\approx70$ pc in radius); (2) 70~$\\mu$m flux, normalized by TIR, remains nearly constant with radius ($L_{70}\\approx45\\%~L_{\\mbox{TIR}}$) from $\\sim10$ to 400 pc; (3) 70~$\\mu$m flux, normalized by TIR, has the smallest fractional error ($0.05-0.12$ out to 220 pc); and (4) 70 $\\mu$m flux, normalized by TIR, does not systematically differ between the LMC and SMC.  Radial SFR density plots, using the obscured indicator (TIR) and unobscured indicators (UV, {\\Ha}), are computed using the SFR recipes of \\citet{kenn98}.  The computed SFRs for all indicators peak near the center of each {\\HII} region.  The SFRs for the LMC are larger than for the SMC due to the larger luminosities of the LMC {\\HII} regions. Within the limits of our resolution, the spatial similarities between UV, {\\Ha}, and 24~$\\mu$m are similar to results in the literature.  A modified version of the obscured SFR equation from \\citet{kenn98} is created that depends on a conversion factor from TIR to a monochromatic IR band.  The conversion factors for the 8, 24, and 160~$\\mu$m bands all have large uncertainties and are dependent on radius and/or host galaxy.  The conversion factor for 70~$\\mu$m is constant, within the low uncertainties, at distances from 10 to 300 pc, and between the LMC and SMC.  We produce a final modified obscured SFR equation using a single averaged 70~$\\mu$m conversion constant. This \\citet{kenn98} modified equation is applicable for Magellanic-like {\\HII} regions and for aperture sizes of $10-300$ pc radius.  The LMC and SMC have different environmental properties than those galaxies of earlier Hubble types.  The dust emission observed around Magellanic Cloud {\\HII} regions may not be indicative of dust emission around {\\HII} regions in other galaxies.  The analysis in this work would be greatly benefited by including {\\HII} regions in a larger sample of galaxies.  By analyzing the 30~Doradus IR properties, we found possible caveats in using 70~$\\mu$m as an obscured star formation indicator.  If {\\HII} regions in a particular galaxy have, on average, similar fluxes as 30~Doradus, then the 70~$\\mu$m emission, relative to the TIR, will exhibit more scatter.  However, we do not expect this to be the case because 30~Doradus is on the extreme high end of {\\HII} region {\\Ha}-derived luminosity functions \\citep{kenn89}.  Another general caveat to mention is that this work applies directly to {\\HII} regions and their immediate surroundings.  It is general practice for those studying SFRs over entire galaxies to use the SFR recipes calculated from {\\HII} regions, despite the added uncertainties of going from a local star forming region to an entire galaxy.  The results from this work, applied to entire galaxies, will be subject to similar uncertainties.  The recent launch of the \\textit{HERSCHEL} telescope will greatly increase our understanding of star formation in galaxies.  At 3.5~m, not only is \\textit{HERSCHEL} the largest space telescope ever launched, but the PACS photometer, with the ability to observe at 75~$\\mu$m, 110~$\\mu$m, or 175~$\\mu$m, will provide a measure of {\\HII} region SEDs at wavelengths near their peak emissions.  The SPIRE camera will observe at 250~$\\mu$m, 350~$\\mu$m, and 500~$\\mu$m (albeit with a lower resolution than PACS) and can provide a better constraint on the Rayleigh--Jeans side of the dust blackbody emission.  NASA's 2.5~m \\textit{SOPHIA} telescope will also have a suite of high-resolution instruments capable of observing star forming regions throughout the IR, including near the expected peaks of the {\\HII} region SEDs.  This work is based on observations made with the \\textit{Spitzer Space Telescope}, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under NASA contract 1407.  The SAGE-LMC project has been supported by NASA/Spitzer grant 1275598 and Meixner's efforts have had additional support from NASA NAG5-12595."
    },
    "1005/1005.2519_arXiv.txt": {
        "abstract": "{The Polaris Flare cloud region contains a great deal of extended emission. It is at high declination and high Galactic latitude. It was previously seen strongly in IRAS Cirrus emission at 100 microns. We have detected it with both PACS and SPIRE on {\\it Herschel}. We see filamentary and low-level structure. We identify the five densest cores within this structure. We present the results of a temperature, mass and density analysis of these cores. We compare their observed masses to their virial masses, and see that in all cases the observed masses lie close to the lower end of the range of estimated virial masses. Therefore, we cannot say whether they are gravitationally bound prestellar cores. Nevertheless, these are the best candidates to be potential prestellar cores in the Polaris cloud region.} ",
        "introduction": "In this paper we present observations, performed with the ESA {\\it Herschel} Space Observatory (Pilbratt et al 2010), of the Polaris Flare region. In particular we use the large collecting area and powerful science payload of {\\it Herschel} to perform imaging photometry using the PACS (Poglitsch et al 2010) and SPIRE (Griffin et al 2010) instruments. These observations were carried out as part of the guaranteed-time key programme to map most of the Gould Belt star-forming regions with {\\it Herschel} (Andr\\'e et al 2010). The Polaris Flare was first detected in HI as a spur of gas that appears to rise more than 30$^\\circ$ out of the Galactic Plane. This region is an area rich in IRAS cirrus emission (e.g. Low et al 1984), and is sometimes known as the Polaris Cirrus Cloud. It was mapped in CO by Heithausen \\& Thaddeus (1990). On the large scale this cloud appears to merge with the Cepheus Flare cloud (e.g. Kirk et al 2009), and both clouds extend to high Galactic latitude.  \\begin{figure*} \\centering{\\includegraphics[width=\\textwidth]{14618fg1.eps}} \\caption{\\label{5bands}The densest part of the Polaris Flare region at some of the observed wavebands. Upper row: 160$\\mu$m from PACS, and 250$\\mu$m and 350$\\mu$m from SPIRE. Lower row: False-colour image (where 160$\\mu$m is shown in blue, 250$\\mu$m is shown in green, and 350$\\mu$m is shown in red), column density map (where red is $<$4, blue is 4--8, and yellow is $>$8 $\\times$ 10$^{21}$ cm$^{-2}$), and colour temperature map (where blue is 10--11~K and yellow is 12--13~K). The contour levels on the column density map start  at 4 $\\times$ 10$^{21}$ cm$^{-2}$, and the interval between successive contours is 1.5 $\\times$ 10$^{21}$ cm$^{-2}$. The same contours are repeated on the temperature map for ease of location. Five sources are seen above a column density of 4 $\\times$ 10$^{21}$ cm$^{-2}$. These are labelled cores 1--5 (in order of increasing R.A.) on the last two panels and are discussed in the text. The loop (loop 1) discussed in the text (containing cores 4 \\& 5) is clearly visible in all images. The reddest features on the false-colour image are the coldest, and the loop shows up clearly as redder than the surroundings. Likewise in the temperature map, the loop shows up as blue, indicating that it is the coldest feature on the map. The position of the IRAS source (IRAS~01432+8725) is marked with a star on the last two panels (adjacent to core 4).} \\end{figure*}  One of the denser regions in the cloud is known as molecular cloud 123.5+24.9, or MCLD~123.5+24.9 (e.g. Bensch et al 2003) -- hereafter MCLD~123 -- at a distance of 150~pc (Bensch et al 2003). It shows strong extended IRAS 100-$\\mu$m emission and is generally believed to be gravitationally unbound with a mass of $\\sim$18--32~M$_\\odot$ (Grossmann et al 1990; Bensch et al 2003). A CO study by Falgarone et al (1998) revealed a curved filament in MCLD~123 in $^{13}$CO and C$^{18}$O -- both in the J=2--1 transition. This filament is also apparent in some narrow velocity channels in the same transition of $^{12}$CO (Falgarone et al 1998).  There is one IRAS source in the region, IRAS~01432+8725. This is listed in the IRAS catalogue as having a flux density at 100~$\\mu$m of 2.88~Jy, but only upper limits at the other IRAS wavebands. There is also one {\\it Spitzer} source that was only detected at a wavelength of 24~$\\mu$m at coordinates R.A. (2000) $=$ 01$^{\\rm h}$ 58$^{\\rm m}$ 27.5$^{\\rm s}$, Dec. (2000) $=$ $+$87\\degr\\ 40\\arcmin\\ 07\\arcsec\\ . It has a peak flux density at 24~$\\mu$m of 1.3~mJy/beam, where the {\\it Spitzer} beam at this wavelength is 7~arcsec. This detection lies in a {\\it Spitzer} calibration field in an unpublished archival dataset (AOR~33136386). ",
        "conclusions": "We have presented {\\it Herschel} data of the Polaris Flare dark cloud region, and in particular the region MCLD~123. We found a great deal of extended emission at wavelengths from 70 to 500~$\\mu$m with both PACS and SPIRE. We noted some filamentary and low-level structure. We identified the five densest cores within this structure. We carried out a temperature, mass and density analysis of the cores. We compared their observed masses to their virial masses, and found that the observed masses are on the lower limit of the range of their estimated virial masses, and thus we cannot say for certain whether they are gravitationally bound."
    },
    "1005/1005.2175_arXiv.txt": {
        "abstract": "Certain oscillatory features in the primordial scalar power spectrum are known to provide a better fit to the outliers in the cosmic microwave background data near the multipole moments of $\\ell=22$ and $40$. These features are usually generated by introducing a step in the popular, quadratic potential describing the canonical scalar field. Such a model will be ruled out, if the tensors remain undetected at a level corresponding to a tensor-to-scalar ratio of, say, $r\\simeq 0.1$. In this work, in addition to the popular quadratic potential, we investigate the effects of the step in a small field model and a tachyon model. With possible applications to future datasets (such as PLANCK) in mind, we evaluate the tensor power spectrum exactly, and include its contribution in our analysis. We compare the models with the WMAP (five as well as seven-year), the QUaD and the ACBAR data. As expected, a step at a particular location and of a suitable magnitude and width is found to improve the fit to the outliers (near $\\ell=22$ and $40$) in all these cases. We point out that, if the tensors prove to be small (say, $r\\lesssim 0.01$), the quadratic potential and the tachyon model will cease to be viable, and more attention will need to be paid to examples such as the small field models. ",
        "introduction": "Inflation continues to remain the most promising paradigm for describing the origin of the perturbations in the early universe. It has been performing remarkably well against the observational data, and the challenge before the other competing scenarios is to match the simplicity and efficiency of inflation. Many models of inflation lead to an epoch of slow roll lasting for, say, $50$-$60$ e-folds, as is required to resolve the horizon problem. It is well known that slow roll inflation leads to a featureless, nearly scale invariant, power law, primordial scalar spectrum. Such a spectrum, along with the assumption of a spatially flat, concordant $\\Lambda$CDM background cosmological model, provides a good fit to the recent observations of the anisotropies in the Cosmic Microwave Background (CMB) by different missions such as the Wilkinson Microwave Anisotropy Probe (WMAP)~\\cite{wmap-5,wmap-7}, the QUEST at DASI (QUaD)~\\cite{quad-2009}, and the Arcminute Cosmology Bolometer Array Receiver (ACBAR)~\\cite{acbar-2008}.  The efficacy of the inflationary scenario also seems to be responsible for an important drawback. Though, as a paradigm, inflation can be considered to be a success, it would be fair to say that we are rather far from converging on a specific model or even a class of models of inflation. There exist a plethora of inflationary models that remain consistent with the data. We mentioned above that a nearly scale invariant, power law, scalar spectrum fits the observations of the anisotropies in the CMB quite well. However, there exist a few data points at the lower multipoles---notably, at the quadrapole ($\\ell=2$) and near the multipole moments of $\\ell=22$ and $40$---which lie outside the cosmic variance associated with the power law primordial spectrum. Needless to add, statistically, a few outliers in a thousand or so data points can always be expected. These outliers were noticed in the WMAP first-year data, and they continue to be present even in the most recent, seven-year data, making them unlikely to be artifacts of data analysis. It is possible that they actually indicate certain non-trivial inflationary dynamics. In that case, these outliers are important from the phenomenological perspective of attempting to constrain the models from the data, because only a more restricted class of inflationary models can be expected to provide an improved fit to these outliers. Therefore, it is a worthwhile exercise to systematically explore models that lead to specific deviations from the standard power law, inflationary perturbation spectrum, and also provide an improved fit to the data.  Various efforts towards a model independent reconstruction of the primordial spectrum from the observed pattern of the CMB anisotropies seem to indicate the presence of certain features in the spectrum~\\cite{rc}. (However, we should add that there also exist other views on the possibility of features in the primordial spectrum; in this context, see, for example, Refs.~\\cite{rc-wf}.) In particular, a burst of oscillations in the primordial spectrum seems to provide a better fit to the CMB angular power spectrum near the multipole moments of $\\ell=22$ and $40$. Generating these oscillations requires a short period of deviation from slow roll inflation~\\cite{starobinsky-1992,dvorkin-2010}, and such a departure has often been achieved by introducing a small step in the popular, quadratic potential describing the canonical scalar field (see Refs.~\\cite{adams-2001,wmap-1,covi-2006-2007,mortonson-2009}; for a discussion on other models, see, for instance, Refs.~\\cite{joy-2008-2009,jain-2009}). At the cost of three additional parameters which characterize the location, the height and the width of the step, it has been found that this model provides a considerably better fit to the CMB data with the least squares parameter $\\chi_{\\rm eff}^{2}$ typically improving by about $7$, when compared to the nearly scale invariant spectrum that would have resulted in the absence of the step~\\cite{covi-2006-2007,mortonson-2009}. But, such a chaotic inflation model leads to a reasonable amount of tensors, and these models will be ruled out if tensors are not detected corresponding to a tensor-to-scalar ratio of, say, $r\\simeq 0.1$.  Our aims in this paper are twofold. Firstly, we wish to examine whether, with the introduction of a step, other inflationary models too perform equally well against the CMB data, as the quadratic potential does. Secondly, we would also like to consider a model that leads to a tensor-to-scalar ratio of $r<0.1$, so that suitable alternative models exist if the tensor contribution turns out to be smaller. Motivated by these considerations, apart from revisiting the popular quadratic potential, we shall investigate the effects of the step in a small field model~(in this context, see, for example, Ref.~\\cite{efstathiou-2006}) and a tachyon model~\\cite{steer-2004}. Also, with possible applications to future datasets in mind (such as the ongoing PLANCK mission~\\cite{planck}), we shall evaluate the tensor power spectrum exactly, and include its contribution in our analysis. We shall compare the models with the CMB data from the WMAP, the QUaD and the ACBAR missions. We shall consider the five as well as the seven-year WMAP data~\\cite{wmap-5,wmap-7}, the QUaD June 2009 data~\\cite{quad-2009} and the ACBAR 2008 data~\\cite{acbar-2008} to arrive at the observational constraints on the inflationary parameters. We find that, as one may expect, a step at a suitable location and of a certain magnitude and width improves the fit to the outliers (near $\\ell=22$ and $40$) in all the cases. We point out that, if the amplitude of the tensors prove to be small, the quadratic potential and the tachyon model will become inviable, and we will have to turn our attention to examples such as the small field models.  The remainder of this paper is organized as follows. In the following section, we shall outline the different inflationary models that we shall be considering in this work. In Section~\\ref{sec:method}, we shall discuss the methodology that we adopt for comparing the inflationary models with the data, the datasets that we use for our analysis, and the priors on the various parameters that we work with. In Section~\\ref{sec:results}, we shall present the results of our comparison of the theoretical CMB angular power spectra that arise from the various models with the WMAP five-year as well as seven-year data, the QUaD and the ACBAR data. We shall tabulate the best fit values that we obtain on the background cosmological parameters and the parameters describing the inflationary models. We shall also illustrate the constraints that we arrive at on the parameters describing the step in the case of the small field model. Further, we shall explicitly show that the models with the step perform better against the data because of the fact that they lead to an improvement in the fit to the outliers around $\\ell=22$ and $40$. In Section~\\ref{sec:sps}, we shall illustrate the scalar power spectra and the CMB angular power spectra corresponding to the best fit values of the parameters of some of the models that we consider. Finally, in Section~\\ref{sec:summary}, we shall close with a brief summary, and a few comments on certain implications of our results.  Note that we shall work in units such that $\\hbar=c=(8\\,\\pi\\,G)=1$. Moreover, we shall assume the background cosmological model to be the standard, spatially flat, $\\Lambda$CDM model. ",
        "conclusions": "\\label{sec:summary}  In this work, we have investigated the effects of introducing a step in certain inflationary models. In addition to revisiting the case of the quadratic potential that has been considered earlier, we have studied the effects of the step in a small field model and a tachyon model. The introduction of the step leads to a small deviation from slow roll inflation, which results in a burst of oscillations in the scalar power spectrum. These oscillations, in turn, leave their imprints as specific features in the CMB angular power spectrum. Actually, we have also evaluated the tensor power spectrum exactly, and have included it in our analysis. We believe that this is a timely effort considering the fact that results from, say, the ongoing PLANCK mission might necessitate such an analysis. Upon comparing the inflationary models with the WMAP, the QUaD and the ACBAR data, we find that, with the step, all the models lead to an improvement in $\\chi_{\\rm eff}^{2}$ by about $7$-$9$ over the smooth, nearly scale invariant, slow roll spectrum, at the expense of three additional parameters describing the location, the height and the width of the step in the inflaton potential. The output of the WMAP likelihood code and a plot of the difference in $\\chi_{\\rm eff}^{2}$ with and without the step clearly illustrate that the improvement occurs because of a better fit to the data at the lower multipoles due to the presence of the step. Evidently, if future observations indicate that the amplitude of the tensors are rather small, then the quadratic potential and the tachyon model will be ruled out, while a suitable small field model with a step can be expected to perform well against the data.  The introduction of the step in an inflationary model can possibly be viewed as an abrupt change in a potential parameter~\\cite{adams-2001}. But, it has to be admitted that it is rather ad-hoc, and one needs to explore the generation of features and a resulting improvement in the fit in better motivated inflationary models. Two field models offer such a possibility. For instance, with suitably chosen parameters, the two field models can easily lead to a brief departure from slow roll inflation (in this context, see, for example, Refs.~\\cite{joy-2008-2009,cline-2003,hunt-2004-2007}). However, iso-curvature perturbations arise whenever more than one field is involved~\\cite{gordon-2000,tsujikawa-2003-2005}, and they need to be carefully taken into account when comparing these models with the data. It will be a worthwhile effort to systematically explore the two field models, including the effects due to the iso-curvature perturbations, in an attempt to fit the outliers near the multipole moments of $\\ell=22$ and $40$.  Over the last few years, it has been recognized that primordial non-Gaussianity can act as a powerful observational tool that can help us discriminate further between the various inflationary models. For example, it has been shown that slow roll inflation driven by the canonical scalar fields leads only to a small amount of non-Gaussianity~\\cite{maldacena-2003}. However, recent analysis of the CMB data seem to suggest that the extent of primordial non-Gaussianity may possibly be large (see, for instance, Refs.~\\cite{wmap-7,smith-2009}). It is known that models which lead to features, such as the ones we have considered here, also generate a reasonably large non-Gaussianity (see, for example, Refs.~\\cite{chen-2007-2008}). While the different models that we have considered in this work lead to virtually the same scalar power spectrum and almost the same extent of improvement in the fit (i.e. with the introduction of the step) to the CMB data, it is important to examine whether they lead to the same extent of non-Gaussianity as well. We are currently investigating such issues."
    },
    "1005/1005.2944_arXiv.txt": {
        "abstract": "{ We present NLTE Li abundances for 88 stars in the metallicity range -3.5 $<$ [Fe/H]  $<$ -1.0. The effective temperatures are based on the infrared flux method with improved $E(B-V)$ values obtained mostly from interstellar \\ion{Na}{i}\\,D lines. The Li abundances were derived through MARCS models and high-quality UVES+VLT, HIRES+Keck and FIES+NOT spectra, and complemented with reliable equivalent widths from the literature. The less-depleted stars with [Fe/H]$<-2.5$ and [Fe/H] $>-2.5$ fall into two well-defined plateaus of $A_{\\rm Li}$ $= 2.18 \\; (\\sigma =0.04)$ and $A_{\\rm Li}$  $=2.27 \\; (\\sigma =0.05)$, respectively. We show that the two plateaus are flat, unlike previous claims for a steep monotonic decrease in Li abundances with decreasing metallicities. At all metallicities we uncover a fine-structure in the Li abundances of Spite plateau stars, which we trace to Li depletion that depends on both metallicity and mass. Models including atomic diffusion and turbulent mixing seem to reproduce the observed Li depletion assuming a primordial Li abundance $A_{\\rm Li}$  = 2.64, which agrees well with current predictions ($A_{\\rm Li}$ = 2.72) from standard Big Bang nucleosynthesis. Adopting the Kurucz overshooting model atmospheres increases the Li abundance by +0.08 dex to $A_{\\rm Li}$ = 2.72, which perfectly agrees with BBN+WMAP. } ",
        "introduction": "One of the most important discoveries in the study of the chemical composition of stars was made in 1982 by Monique and Fran\\c cois Spite, who found an essentially constant Li abundance in warm metal-poor stars (Spite \\& Spite 1982), a result interpreted as a relic of primordial nucleosynthesis. Due to its cosmological significance, there have been many studies devoted to Li in metal-poor field stars (e.g., Ryan et al. 1999; Mel\\'endez \\& Ram{\\'{\\i}}rez 2004, hereafter MR04; Boesgaard et al. 2005, hereafter B05; Charbonnel \\& Primas 2005; Asplund et al. 2006, hereafter A06; Shi et al. 2007, hereafter S07; Bonifacio et al. 2007, hereafter B07; Hosford et al. 2009; Aoki et al. 2009; Sbordone et al. 2010, hereafter S10), with observed Li abundances at the lowest metallicities ([Fe/H] $\\sim-$3) from $A_{\\rm Li}$  = 1.94 (B07) to $A_{\\rm Li}$  = 2.37 (MR04).  A primordial Li abundance of $A_{\\rm Li}$  = 2.72$_{-0.06}^{+0.05}$ is predicted (Cyburt et al. 2008; see also Steigman 2009; Coc \\& Vangioni 2010) with the theory  of big bang nucleosynthesis (BBN) and the baryon density obtained from WMAP data (Dunkley et al. 2009), which is a factor of 2-6 times higher than the Li abundance inferred from halo stars. There have been many theoretical studies on non-standard BBN trying to explain the cosmological Li discrepancy by exploring the frontiers of new physics (e.g. Coc et al. 2009; Iocco et al. 2009; Jedamzik \\& Pospelov 2009; Kohri \\& Santoso 2009). Alternatively, the Li problem could be explained by a reduction of the original Li stellar abundance due to internal processes (i.e., by stellar depletion). In particular, stellar models including atomic diffusion and mixing can deplete a significant fraction of the initial Li content (Richard et al. 2005; Piau 2008), although such models depend on largely unconstrained free parameters. On the other hand, it is not easy to reconcile the lack of observed abundance scatter in the Spite plateau with substantial Li depletion (e.g. Ryan et al. 1999; A06). Due to the uncertainties in the Li abundances and to the small samples available, only limited comparisons of models of Li depletion with stars in a broad range of mass and metallicities  have been performed (e.g., Pinsonneault et al. 2002; B05).  The observed scatter in derived Li abundances in previous studies of metal-poor stars can be as low as 0.03 dex (e.g. Ryan et al. 1999; A06), fully consistent with the expected observational errors. Yet, for faint stars observational errors as high as 0.2 dex have been reported (e.g., Aoki et al. 2009). In order to provide meaningful comparisons with stellar depletion models, precise Li abundances for a large sample of stars  are needed. Here we present such a study for the first time for  a large sample of metal-poor stars (-3.5 $<$ [Fe/H] $<$ -1.0) with masses in a relatively broad mass range (0.6-0.9 M$_\\odot$). ",
        "conclusions": "\\subsection{The \\teff cutoff of the Spite plateau}  Despite the fact that Li depletion depends on mass (e.g. Pinsonneault et al. 1992), this variable has been ignored by most previous studies. Usually a cutoff in \\teff is imposed to exclude severely Li-depleted stars in the Spite plateau, with a wide range of adopted cutoffs, such as 5500 K (Spite \\& Spite 1982), 5700 K (B05), 6000 K (MR04; S07) and $\\sim$6100 K for stars with [Fe/H] $< -2.5$ (Hosford et al. 2009).  At a given mass, the \\teff of metal-poor stars has a strong metallicity-dependence (e.g. Demarque et al. 2004). As shown in Figs. 11-12 of M06, the \\teff of turnoff stars increases for decreasing metallicities. Hence, a metallicity-independent cutoff in \\teff may be an inadequate way to exclude low-mass Li-depleted stars from the Spite plateau. As show in Fig. 2, where $A_{\\rm Li}$ in different metallicity bins is shown as a function of \\tsin, stars with lower \\teff in a given metallicity regime are typically the stars with the lowest Li abundances, an effect that can be seen even in the sample stars with the lowest metallicities ([Fe/H] $\\sim -3$). This is ultimately so because the coolest stars are typically the least massive, and therefore have been more depleted in Li (see Sect. 4.3).  In Fig. 3 we show the Li abundance for cutoffs = 5700 K (open circles), 6100 K (filled squares) and 6350 K (filled triangles). Using a hotter cutoff is useful to eliminate the most Li-depleted stars at low metallicities, but it removes from the Spite plateau stars with [Fe/H] $>$ -2. Imposing a hotter \\teff cutoff at low metallicities and a cooler cutoff at high metallicities eliminates the most Li-depleted stars at low metallicities, but keeps the most metal-rich stars in the Spite plateau. We propose such a metallicity-dependent cutoff below.  \\subsection{Two flat Spite plateaus}  Giving the shortcomings of a constant \\teff cutoff, we propose an empirical cutoff of \\teff $=$ 5850 - 180$\\times$[Fe/H]. The stars above this cutoff are shown as stars in the bottom panel of Fig. 3. Our empirical cutoff excludes only the most severely Li-depleted stars, i.e., the stars that remain in the Spite plateau may still be affected by depletion. The less Li-depleted stars in the bottom panel of Fig. 3 show two well-defined groups separated at [Fe/H] $\\sim-2.5$ (as shown below, this break represents a real discontinuity), which have essentially zero slopes (within the error bars) and very low star-to-star scatter in their Li abundances. The first group has $-2.5 \\leq$ [Fe/H] $< -1.0$ and $<$A$_{\\rm Li}$$>_1$ = 2.272 ($\\sigma$=0.051) dex and a slope of 0.018$\\pm$0.026, i.e., flat within the uncertainties. The second group is more metal poor ([Fe/H] $< -2.5)$ and has $<$A$_{\\rm Li}$$>_2$ = 2.184 dex ($\\sigma = 0.036$) dex. The slope of this second group is also zero ($-$0.008$\\pm$0.037). Adopting a more conservative exponential cutoff obtained from $Y^2$ isochrones (Demarque et al. 2004), which for a 0.79 M$_\\odot$ star can be fit by \\teff = 6698 $-$2173 $\\times$ $e^{\\rm [Fe/H]/1.021}$, we would also recover a flat Spite plateau, although only stars with [Fe/H] $> -$2.5 are left using this more restrictive cut-off. Thus, the flatness of the Spite plateau is independent of applying a linear or an exponential cutoff.  Adopting a constant cutoff in \\teff we also find flat plateaus. For example adopting a cutoff of \\teff $>$ 6100 K (filled squares in Fig. 3) we find in the most metal-rich plateau ([Fe/H] $\\geq -2.5$) no trend between Li and [Fe/H] (slope = 0.019$\\pm$0.025, Spearman rank correlation coefficient $r_{\\rm Spearman}$=0.1 and a probability of 0.48 (i.e., 48\\%) of a correlation arising by pure chance for [Fe/H] $\\geq -$2.5), while for the most metal-poor plateau ([Fe/H] $< -$2.5) we also do not find any trend within the errors (slope = 0.058$\\pm$0.072, $r_{\\rm Spearman}$=0.2 and 41\\% probability of a spurious correlation). Using a hotter cutoff (\\teff $>$ 6350 K, filled triangles in Fig. 3) we obtain also two flat plateaus with slope = -0.040$\\pm$0.063 ($r_{\\rm Spearman}$ = -0.2, probability = 60\\%) for [Fe/H] $\\geq -$2.5 and slope = 0.008$\\pm$0.035 ($r_{\\rm Spearman}$ = 0.1, probability = 68\\%) for [Fe/H] $< -$2.5.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{breakli.eps}} \\caption{(a) filled circles show the slope/$\\sigma$ of $A_{\\rm Li}$ vs. [Fe/H] in the range X$_{\\rm min}$ $<$ [Fe/H] $<-1.0$, where X$_{\\rm min}$ is in the interval [-3.5,-1.25]. For all values of X$_{\\rm min}$ $\\geq -2.5$, the slope is insignificant ($<<$ 3$\\sigma$); only when stars with [Fe/H] $< -2.5$ are included, a slope is forced in $A_{\\rm Li}$ vs. [Fe/H]. Open circles show the slope/$\\sigma$ for the range $-3.25<$ [Fe/H] $<$ X$_{\\rm max}$, where $-$3.0 $\\leq$ X$_{\\rm max}$ $\\leq -$1.0. For X$_{\\rm max}$ $< -2.5$ the slope is negligible ($<<$ 3$\\sigma$), but when stars with [Fe/H] $\\geq -2.5$ are included, a slope is generated. (b) The correlation coefficient $r_{\\rm Spearman}$ is shown by triangles, with filled and open symbols with similar meaning as in panel (a). The probabilities of a correlation between $A_{\\rm Li}$ and [Fe/H] by pure chance, associated to the filled and open triangles, are shown by solid and dashed lines, respectively. No meaningful correlation (probability $<<$ 1) exists for either [Fe/H] $< -$2.5 or [Fe/H] $\\geq -$2.5. } \\label{libreak} \\end{figure}   In Fig. 4 we demonstrate that the break at [Fe/H] $\\sim -2.5$ is statistically significant. In panel (a) the slope/$\\sigma$ in the $A_{\\rm Li}$ vs. [Fe/H] plot for the range X$_{\\rm min}$ $<$ [Fe/H] $<-1.0$  are shown as filled circles, where X$_{\\rm min}$ varies within [-3.50,-1.25]. For X$_{\\rm min}$ $\\geq -2.5$, the slope is insignificant ($<<$ 3$\\sigma$), and only when stars with [Fe/H] $< -2.5$ are included a measurable slope is forced in the  $A_{\\rm Li}$ vs. [Fe/H] relation. The opposite test is shown by open circles, where we show slope/$\\sigma$ for the range $-$3.25 $<$ [Fe/H] $<$ X$_{\\rm max}$, where X$_{\\rm max}$ changes from [$-$3.00,$-$1.00]. For X$_{\\rm max}$ $< -2.5$ the slope is negligible ($<<$ 3$\\sigma$), and only when stars with [Fe/H] $\\geq -2.5$ are included, a slope is produced between $A_{\\rm Li}$ and [Fe/H]. The correlation coefficient $r_{\\rm Spearman}$ and the probability of a correlation between $A_{\\rm Li}$ and [Fe/H] by pure chance are shown in panel (b). Again, this plot shows that no correlation between $A_{\\rm Li}$ and [Fe/H] exists {\\it within} the two groups ($-3.5<$ [Fe/H] $<-2.5$ and $-2.5\\geq$ [Fe/H] $\\geq-1.0$), and that only when stars {\\it between} the two groups are mixed, significant (probability $\\sim$ 0) correlations of $A_{\\rm Li}$ with [Fe/H] are generated. Systematically lower E(B-V) values (by $\\sim$0.03 mags) in the most metal-poor plateau ([Fe/H]  $< -$2.5) could produce a more Li-depleted plateau. However, those E(B-V) values (Fig.~1) are actually slightly higher ($\\sim$0.007 mags, i.e., $\\sim$33 K) than those for [Fe/H]  $\\geq -$2.5 (mainly due to a high number of unreddened nearby more metal-rich stars), thus not explaining the existence of two different plateaus.  Previous claims of a steep monotonic decrease in Li abundance with decreasing metallicity (e.g. Ryan et al. 1999; A06; B07; Hosford et al. 2009) are probably due to the mix of stars from the two different groups, forcing a monotonic dependence with metallicity. Our large sample of homogeneous and precise Li abundances that covers a broad metallicity range ($-3.5 <$ [Fe/H] $< -1.0$) does not support these claims. Nevertheless, a hint of two different groups in the Spite plateau was already found by A06, who found a change in the slope of the Spite plateau at [Fe/H] $\\approx$ $-$2.2. Also, in the combined A06+B07 sample (Fig. 7 of B07), there are two different groups: stars with [Fe/H] $\\gtrsim -2.6$ have $A_{\\rm Li}$ $>$ 2.2, while stars with [Fe/H] $\\lesssim -2.6$ have $A_{\\rm Li}$ $<$ 2.2. Although in the study by MR04 a flat Spite plateau is found in the range $-3.4 <$ [Fe/H] $< -1$, this is due to the overestimation of \\teff below [Fe/H] $< -2.5$, thus overestimating $A_{\\rm Li}$ at low metallicities and forcing a flat plateau from [Fe/H] = $-3.4$ to $-$1.   \\subsection{Correlation between Li and mass}  Models of Li depletion predict that the least massive stars are the most depleted in Li, but due to the limitations of previous samples, these predictions have not been thoroughly tested at different metallicity regimes in metal-poor stars.\\footnote{Except for the work of B05, who provide comparisons at different metallicities but for cool severely Li-depleted dwarfs, i.e. probably of lower mass than most stars shown in Fig. 5} In Fig. 5 we show our Li abundances as a function of stellar mass for different metallicity ranges. As can be seen, the Li plateau stars have a clear dependence with mass for all metallicity regimes. Excluding stars with mass $<$ 0.7M$_\\odot$ (including those stars will result in even stronger correlations), linear fits result in slopes of 6, 3, 2, 2 dex M$_\\odot^{-1}$ for stars in the metallicity ranges $-1.3 <$ [Fe/H] $< -1.0$, $-1.6 <$ [Fe/H] $< -1.3$, $-2.5 <$ [Fe/H] $< -1.6$, and $-3.5 <$ [Fe/H] $< -2.5$. The slopes are significant at the 8, 2, 5, 1 $\\sigma$ level, respectively. The correlation coefficient $r_{\\rm Spearman}$ is 0.9, 0.6, 0.6, 0.3, and the probability of a correlation between  $A_{\\rm Li}$ and mass arising by pure chance is very small: 5$\\times 10^{-5}$, 3$\\times 10^{-2}$, 1$\\times 10^{-3}$, and 1.3$\\times 10^{-1}$, for stars in the same metallicity ranges as above. Thus, the correlations of $A_{\\rm Li}$ and mass in different metallicity regimes are very significant.  Recently, Gonz\\'alez Hern\\'andez et al. (2008, hereafter G08) have studied the metal-poor ([Fe/H] $\\sim-$3.5) double-lined spectroscopic binary BPS CS22876-032, providing thus crucial data to test our Li-mass trend. With our method and the stellar parameters of G08, we obtain a mass ratio of 0.89, very close to their value (0.911$\\pm$0.022) obtained from an orbital solution. For the primary we obtain $M_A$ = 0.776 M$_\\odot$, and adopting the mass ratio of G08, $M_B$ = 0.707 M$_\\odot$ is obtained for the secondary. The LTE Li abundances were taken from G08 and corrected for NLTE effects ($\\sim-$0.05 dex). The components of the binary are shown as circles in Fig. 5, nicely following the trend of the most metal-poor stars. Including this binary in our sample would strengthen the Li-mass correlation of stars with [Fe/H] $< -$2.5. A slope = 3 dex  M$_\\odot^{-1}$ significant at the 3$\\sigma$ level is obtained, with $r_{\\rm Spearman}$ = 0.5 and a low probability (2$\\times 10^{-2}$) of the trend being spurious.  While mass-dependent Li depletion is expected from standard models of stellar evolution, this should only occur at significantly lower masses than considered here. These stellar models only predict very minor $^7$Li depletion ($\\la 0.02$\\,dex) for metal-poor turn-off stars (e.g. Pinsonneault et al. 1992), which is far from sufficient to explain the $\\sim$0.5\\,dex discrepancy between the observed Li abundance and predictions from BBN+WMAP. Bridging this gap would thus require invoking stellar models that include additional processes normally not accounted for, such as rotationally-induced mixing or diffusion. In Fig. 5 we confront the predictions of Richard et al. (2005) with our inferred stellar masses and Li abundances. The models include the effects of atomic diffusion, radiative acceleration, and gravitational settling, but moderated by a parameterized turbulent mixing (T6.0, T6.09, and T6.25, where higher numbers mean higher turbulence); so far only the predictions for [Fe/H]=$-$2.3 are available for different turbulent mixing models. The agreement is very good when adopting a turbulent model of T6.25 (see Richard et al. for the meaning of this notation) and an initial $A_{\\rm Li}  = 2.64$; had the Kurucz convective overshooting models been adopted, the required initial abundance to explain our observational data would correspond to $A_{\\rm Li} = 2.72$. Two other weaker turbulence models that produce smaller overall Li depletions are also shown in Fig. 5, but they are less successful in reproducing the observed Li abundance pattern. Our best-fitting turbulence model (T6.25) is different from that (T6.0) required to explain the abundance pattern in the globular cluster NGC\\,6397 at similar metallicity (Korn et al. 2006; Lind et al. 2009b). Another problem with adopting this high turbulence is that the expected corresponding $^6$Li depletion would amount to $>1.6$\\,dex and thus imply an initial $^6$Li abundance at least as high as the primordial $^7$Li abundance if the claimed $^6$Li detections in some halo stars (A06; AM08) are real.  Our results imply that the Li abundances observed in Li plateau stars have been depleted from their original values and therefore do not represent the primordial Li abundance. It appears that the observed Li abundances in metal-poor stars can be reasonably well reconciled with the predictions from standard Big Bang nucleosynthesis (e.g. Cyburt et al. 2008) by means of stellar evolution models that include Li depletion through diffusion and turbulent mixing (Richard et al. 2005). We caution however that although encouraging, our results should not be viewed as proof of the Richard et al. models until the free parameters required for the stellar modelling are better understood from physical principles.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{masslimodel.eps}} \\caption{Li abundances as a function of stellar mass in different metallicity ranges. The metal-poor ([Fe/H] $\\sim-$3.5) binary BPS CS22876-032 (G08) is represented by open circles, nicely fitting the Li-mass trend. Models at [Fe/H] = $-2.3$ including diffusion and T6.0 (short dashed line), T6.09 (dotted line) and T6.25 (solid line) turbulence (Richard et al. 2005) are shown. The models were rescaled to an initial $A_{\\rm Li}$=2.64 (long dashed line) and by $\\Delta M = +0.05$\\,M$_\\odot$. } \\label{massli} \\end{figure}"
    },
    "1005/1005.0346_arXiv.txt": {
        "abstract": "Using a 3D GCM, we create dynamical model atmospheres of a representative transiting giant exoplanet, HD 209458b. We post-process these atmospheres with an opacity code to obtain transit radius spectra during the primary transit. Using a spectral atmosphere code, we integrate over the face of the planet seen by an observer at various orbital phases and calculate light curves as a function of wavelength and for different photometric bands. The products of this study are generic predictions for the phase variations of a zero-eccentricity giant planet's transit spectrum and of its light curves. We find that for these models the temporal variations in all quantities and the ingress/egress contrasts in the transit radii are small ($< 1.0$\\%). Moreover, we determine that the day/night contrasts and phase shifts of the brightness peaks relative to the ephemeris are functions of photometric band.  The $J$, $H$, and $K$ bands are shifted most, while the IRAC bands are shifted least.  Therefore, we verify that the magnitude of the downwind shift in the planetary ``hot spot\" due to equatorial winds is strongly wavelength-dependent.  The phase and wavelength dependence of light curves, and the associated day/night contrasts, can be used to constrain the circulation regime of irradiated giant planets and to probe different pressure levels of a hot Jupiter atmosphere.  We posit that though our calculations focus on models of HD 209458b similar calculations for other transiting hot Jupiters in low-eccentricity orbits should yield transit spectra and light curves of a similar character. ",
        "introduction": "\\label{intro}    Circa April 2010, more than 70 exoplanets had been discovered transiting their primary stars\\footnote{see J. Schneider's Extrasolar Planet Encyclopaedia at http://exoplanet.eu, the Geneva Search Programme at http://exoplanets.eu, and the Carnegie/California compilation at http://exoplanets.org}.  Transits are useful in the characterization of exoplanets because they break the the mass/orbital-inclination degeneracy of radial-velocity measurements.  Moreover, a planet's radius can be directly measured, yielding mass (\\mp) and radius (\\rp) correlations with which theorists can extract useful information about bulk composition and structure and can attempt to fit radius evolution models (Guillot et al. 1996; Burrows et al. 2000; Brown et al. 2001; Hubbard et al. 2001; Baraffe et al. 2003; Chabrier et al. 2004; Charbonneau et al. 2007).  However, since atmospheric opacity is a function of wavelength, a planet's transit radius is also a function of wavelength (Seager \\& Sasselov 2000; Brown 2001; Spiegel et al. 2007; Fortney et al. 2003,2010).  The variation with wavelength in the measured radius (actually an impact parameter) provides an ersatz spectrum directly related to the planet's atmospheric composition near the terminator.  This radius spectrum, unlike the planet's spectrum at secondary eclipse, is more dependent upon composition than upon the temperature profile, and so provides complementary information to a planet's own direct emissions.  The latter can be measured during secondary eclipse, but also during the traverse by the planet of its orbit as it traces out its phase light curve.  The transit radius spectrum and the planet emission spectrum as a function of phase can, therefore, together help constrain planet properties (Fortney et al. 2006,2010).  However, the stellar irradiation of such an exoplanet severely breaks what would otherwise be quasi-spherical symmetry, producing large day-night contrasts in thermal structure, zonal flows and banding, and violent atmospheric dynamics (Showman \\& Guillot 2002; Guillot \\& Showman 2002; Cho et al. 2003,2008; Burkert et al. 2005; Cooper \\& Showman 2005,2006; Showman et al. 2008,2009; Showman, Cho, \\& Menou 2010; Dobbs-Dixon \\& Lin 2008; Dobbs-Dixon, Cumming, \\& Lin 2010; Rauscher et al. 2007,2008; Menou \\& Rauscher 2009; Rauscher \\& Menou 2010).  Importantly, such behavior may have photometric and spectral signatures. It is the exploration of such signatures that motivates this paper. The state-of-the-art in the 3D modeling of exoplanet atmospheres is still evolving and has not yet reached a level of maturity where detailed predictions for each known exoplanet are robust.  Hence, in this paper we focus on a few generic ideas and conjectures that emerge from our modeling efforts.  The procedure we have pursued is the following:  First, using a 3D general circulation model (GCM) we derive dynamic model atmospheres at various epochs after they have achieved a steady state. Second, we post-process these 3D model atmospheres with a spectral atmosphere code to obtain transit radius spectra at various epochs during ingress, egress, and total transit.  We assume equilibrium molecular compositions and solar metallicity.  Third, using the same 3D model atmosphere, we integrate over the ``visible\" disk and calculate phase light curves as a function of wavelength and for various standard photometric bands.  The results are generic predictions which, though not expected to be quantitatively precise and constraining on any particular giant exoplanet, nevertheless contain qualitatively interesting features that should inform future measurements.  For specificity, we focus on two models of HD 209458b (Charbonneau et al. 2000; Henry et al. 2000), one with and one without an ``extra absorber\" and a thermal inversion (Hubeny, Burrows, \\& Sudarsky 2003; Burrows et al. 2007; Fortney et sl. 2008; Spiegel et al. 2009; Knutson et al. 2008), but suggest our qualitative results are generic beyond this planet and this modeling paradigm (see also Fortney et al. 2006,2010).  In \\S\\ref{techniques}, we summarize our methodology and techniques.  Then, in \\S\\ref{describe} we describe the 3D models and various of their salient characteristics. We go on in \\S\\ref{wave_rad} to present our results for the wavelength-dependent transit radius and ingress-egress asymmetries and in \\S\\ref{phase} we turn to a full discussion of the derived light curves.  This section contains not only wavelength-dependent planet-star flux ratios as a function of phase, but the phase variation of several photometric planet band fluxes.  In \\S\\ref{conclusions}, we summarize our general conclusions. ",
        "conclusions": "\\label{conclusions}   Using a 3D GCM with Newtonian cooling and day- and nightside equilibrium temperature profiles based on two 1D spectral models (with and without a thermal inversion and a hot upper atmosphere), we created dynamical models of the transiting giant exoplanet HD 209458b. We post-processed these 3D model atmospheres with a detailed opacity code to obtain transit radius spectra at various epochs during the primary transit. Then, using a spectral atmosphere code, we integrated over the face of the planet seen by an observer at various orbital phases and calculated light curves as a function of wavelength and for nine different photometric bands from $R$ ($\\sim$0.6 \\mic) through IRAC 4 ($\\sim$8 \\mic). The products of this study are generic predictions for the character of the expected phase variations both of a giant planet's transit spectrum and of its light curves. Since our GCM employed Newtonian cooling (and not radiative transfer using opacities that corresponded with those used in the post-processing), the calculations are slightly inconsistent.  Nevertheless, the results have interesting features that should inform future measurements. We found that the temporal variations in the derived integral quantities for this model suite are small (less than 1\\%) and that the ingress/egress contrasts due to zonal flows, while also small ($< 0.5 - 1.0$ \\%), are most manifest due to the scale-height asymmetries near the terminators and in the water bands. While models we have generated with and without the methane and carbon monoxide bands indicate that ingress/egress differences in the CO and CH$_4$ bands might be diagnostic of both carbon chemistry and abundance asymmetries due to zonal flows near the terminators, the differential effects due to different carbon-species abundances at the different terminators are small (less than a few tenths of a percent).  A similar conclusion was reached by Fortney et al. (2010), who nevertheless found a slightly larger quantitative effect.  Since the transit radius of HD 209458b has been measured in the optical, it should not be allowed to differ when comparing models with and without an upper-atmosphere absorber and the predicted radii at other wavelengths must be determined relative to it.  When this is done, the predicted radii in the near- and mid-infrared are very different (by as much as 0.05 \\rj), though our current models do a poor job of fitting all the transit radius data for HD 209458b simultaneously (e.g., Knutson et al. 2007; Beaulieu et al. 2010; Richardson et al. 2006). The reason for this discrepancy is currently unknown.  We determined that the angular phase shifts of the brightness peaks relative to the orbital ephemeris are functions of photometric band.  The $J$, $H$, and $K$ bands are shifted most, by as much as $\\sim$45$^{\\circ}$, while the IRAC bands are shifted least. This is because different bands are spectrally formed at different pressure depths and the zonal winds that advect heat are depth-dependent. The IRAC band photospheres are generally at altitude, while those in the near infrared are at deeper pressures near and above $\\sim$0.1 bars.  Therefore, the question of the magnitude of the downwind shift in the planetary ``hot spot\" due to equatorial winds is nuanced, depending upon the waveband in which measurements are made.  This wavelength dependence of the phase shift in the brightness light curve, and the associated day/night contrasts, can in principle be used to constrain the circulation regimes of irradiated giant planets and to probe different pressure levels of a hot Jupiter atmosphere.  Though our calculations focused on models of HD 209458b, similar calculations for other transiting hot Jupiters in low-eccentricity orbits should yield transit spectra and light curves of a similar character.  As the subject of comparative exoplanetology matures, and JWST comes online, such observational manifestations of global circulation may well become possible for a wide range of irradiated planets. Our models have been constructed to help prepare the way, however imperfectly, for that era."
    },
    "1005/1005.4639.txt": {
        "abstract": "We investigate whether or not the decadal and multi-decadal climate oscillations have an astronomical origin.  Several global surface temperature records since 1850 and records deduced from the orbits of the planets present  very similar power spectra. Eleven frequencies with period between 5 and 100 years closely correspond in the two records. Among them, large climate oscillations with  peak-to-trough amplitude of about  0.1 $^oC$ and 0.25 $^oC$, and periods of about 20 and 60 years, respectively,  are synchronized  to the orbital periods  of Jupiter and Saturn. Schwabe and Hale solar cycles are also visible in the temperature records. A 9.1-year cycle is synchronized to the Moon's orbital cycles. A phenomenological model based on these astronomical cycles can be used to well reconstruct the temperature oscillations since 1850 and to make partial forecasts for the 21$^{st}$ century.  It is found that at least 60\\% of the global warming observed since 1970   has been induced by the combined effect of the above  natural climate oscillations. The partial forecast indicates that climate may stabilize or cool until 2030-2040. Possible physical mechanisms are qualitatively discussed with an emphasis on the phenomenon of collective synchronization of coupled oscillators.\\\\ \\\\  Please cite this article as: Scafetta, N., Empirical evidence for a celestial origin of the climate oscillations and its implications. \\emph{Journal of Atmospheric and Solar-Terrestrial Physics} (2010), doi:10.1016/j.jastp.2010.04.015 ",
        "introduction": "Milankovic [1941] theorized  that variations in eccentricity, axial tilt, and precession of the  orbit of the Earth determine climate patterns such as  the 100,000 year ice age cycles of the Quaternary glaciation over the last few million years. The variation of the orbital parameters of the Earth is  due to the gravitational perturbations induced by the other planets of the solar system,  primarily Jupiter and Saturn.  Over a much  longer time scale the cosmic-ray flux record well correlates with the warm and ice periods of the Phanerozoic during the last 600 million years:  the cosmic-ray flux oscillations are likely due to the changing galactic environment of the solar system as it crosses the spiral arms of the Milky Way [Shaviv, 2003, 2008; Shaviv and Veizer, 2003; Svensmark, 2007]. Over  millennial and secular time scales several authors have found that variations in total solar irradiance  and variations in solar modulated cosmic-ray flux  well correlate with climate changes: see for example: Eddy, 1976; Hoyt  and  Schatten, 1997; White \\emph{et al.}, 1997; van Loon and Labitzke, 2000; Bond \\emph{et al.}, 2001; Kerr, 2001; Douglass and Clader, 2002; Kirkby, 2007; Scafetta and West, 2005, 2007, 2008; Shaviv, 2008; Eichler \\emph{et al.}, 2009; Soon, 2009; Meehl \\emph{et. al.}, 2009; Scafetta 2009, 2010. Also the annual cycle has an evident astronomical origin.  The above results suggest that the dominant drivers of the climate oscillations have a celestial origin. Therefore, it is legitimate to investigate whether the climate oscillations with a time scale between 1 and 100 years, can be interpreted in astronomical terms too.  Global surface temperature has risen [Brohan \\emph{et al.}, 2006] by about 0.8 $^oC$ and 0.5 $^oC$ since 1900 and 1970, respectively. Humans may have partially contributed to this global warming through greenhouse gas (GHG) emissions [IPCC, 2007]. For instance, the IPCC claims that more than 90\\% of the observed warming since 1900 and practically 100\\% of the observed warming since 1970 have had  an anthropogenic origin (see figure 9.5 in  IPCC, AR4-WG1). The latter conclusion derives merely from the fact that climate models referenced by the IPCC cannot explain the warming occurred since 1970 with any known natural mechanism. Therefore, several scientists have concluded that this warming has been caused by anthropogenic GHG emissions that greatly increased during this same period. This theory is known as the \\emph{anthropogenic global warming theory} (AGWT).  However,  the anthropogenic GHG emissions have increased monotonically since 1850 while the global temperature record did not. Several oscillations are seen in the data since 1850, including a global cooling since 2002: see Figure 1. If these climate oscillations are natural, for example induced by astronomical oscillations, they would determine how climate change should be interpreted [Keenlyside \\emph{et al.}, 2008]. In fact,  during its cooling phase a natural multi-decadal oscillation can hide a global warming caused by human GHG emissions or, alternatively, during its warming phase a natural oscillation can accentuate the warming. If the natural oscillations of the climate are not properly recognized and taken into account, important climate patterns, for example the global warming observed from 1970 to 2000, can be erroneously interpreted. Indeed, part of the 1970-2000 warming could have been induced by a multi-decadal natural cycle during its warming phase that the climate models used by the IPCC have not reproduced.  The IPCC [2007] claims that the climate oscillations are induced by some still poorly understood and modeled internal dynamics of the climate system, such as the ocean dynamics. However, the oscillations of the atmosphere and of the ocean, such as the Pacific Decadal Oscillation (PDO) and the Atlantic Multidecadal Oscillation (AMO), may be induced by complex extraterrestrial periodic forcings that are acting on the climate system in multiple ways. Indeed, the climate system is characterized by interesting cyclical patterns that remind astronomical cycles.   \\begin{figure}[t!] \\includegraphics[angle=-90,width=21pc]{figure1.eps} \\caption{ Top: Global surface temperature anomaly (gray) [Brohan \\emph{et al.}, 2006] against the GISS ModelE average simulation (blak) [Hansen \\emph{et al.}, 2007]. The  figure also shows the quadratic upward trend of the temperature. Bottom:  an eight year moving average smooth of the temperature detrended of its upward quadratic trend. This smooth reveals a quasi-60 year modulation. } \\end{figure}       For example, surface temperature records are characterized by decadal and bi-decadal oscillations which are usually found in good correlation with the (11-year) Schwabe and the (22 year) Hale solar cycles [Hoyt and Schatten, 1997; Scafetta and West, 2005; Scafetta, 2009]. However, longer cycles are of interest herein.  Klyashtorin and Lyubushin [2007] and Klyashtorin et al. [2009]  observed that several centuries of climate records (ice core sample, pine tree samples, sardine and anchovy sediment core samples, global surface temperature records, atmospheric circulation index, length of the day index, fish catching  productivity records, etc.) are characterized by  large 50-70 year and 30-year periodic cycles.  The quasi-60 year periodicity has been also found  in secular  monsoon rainfall records from India, in proxies of monsoon rainfall from Arabian Sea sediments and in rainfall over east China [for example see the following works and their references: Agnihotri \\emph{et al.}, 2002; Sinha A. \\emph{et al.} (2005); Goswami \\emph{et al.}, 2006; Yadava  and Ramesh 2007]. Thus, several records indicate that the climate is characterized by a large quasi-60 year periodicity, plus larger secular climatic cycles and smaller decadal cycles. All these cycles cannot be explained with anthropogenic emissions. Errors in the data, other superimposed patterns (for example, volcano effects and longer and shorter cycles) and some chaotic pattern in the dynamics of these signals may sometimes mask the 60-year cycle.  A multi-secular climatic record that shows a clear quasi-60 year oscillation is depicted in Figure 2: the G. Bulloides abundance variation record found in the Cariaco Basis sediments in the Caribbean sea since 1650 [Black \\emph{et al.}, 1999]. This record is an indicator of the trade wind strength in the tropical Atlantic ocean and of the north Atlantic ocean atmosphere variability. This record  shows five 60-year large cycles. These cycles correlate well with the 60-year modulation of the global temperature observed since 1850 (the correlation is negative). On longer time scales, periods of high G. Bulloides abundance correlate well with periods of reduced solar output (the well-known Maunder, Sp\\\"orer, and Wolf minima), suggesting a solar forcing origin of these cycles [Black \\emph{et al.}, 1999].   \\begin{figure}[t!] \\includegraphics[angle=-90,width=21pc]{figure2.eps} \\caption{ Record of G. Bulloides abundance variations (1-mm intervals)   from  1650 to 1990 A.D. (black line) [Black \\emph{et al.}, 1999]. The gray vertical lines highlight 60-year intervals. Five quasi-60 year cycles are seen in this record, which is a proxy for the Atlantic variability since 1650.} \\end{figure}   Patterson et al. (2004) found 60-62 year cycles in sediments and cosmogenic nuclide records in the NE Pacific. Komitov (2009) found similar cycles in the number of the middle latitude auroras from 1700 to 1900. A cycle of about 60 years has been detected in the number of historically recorded meteorite falls in China from AD 619 to 1943 and in the number of witnessed falls in the world  from 1800 to 1974 [Yu \\emph{et al.}, 1983]. Ogurtsov \\emph{et al.} [2002] found   a 60-64 year cycle in $^{10}Be$, $^{14}C$ and Wolf number over the past 1000 years.  The existence of a 60-year signal has been found in the Earth's angular velocity and in the geomagnetic field  [Roberts \\emph{et al.}, 2007]. These results clearly suggest an astronomical origin of the  60-year variability found in several climatic records.  Interestingly,  the traditional Chinese calendar, whose origins can be traced as far back as the 14$^{th}$ century BCE, is arranged in major 60-year cycles [Aslaksen, 1999]. Each year is assigned a name consisting of two components. The first component is one of the 10 \\emph{Heavenly Stems}  (\\emph{Jia, Yi, Bing}, etc.), while the second component is one of the 12 \\emph{Earthly Branches}  that features the names of 12 animals (\\emph{Zi, Chou, Yin}, etc.). Every 60 years the stem-branch cycle repeats. Perhaps, this sexagenary cyclical calendar was inspired by climatic and astronomical observations.     Some studies [Jose, 1965; Landscheidt, 1988, 1999; Charv\\'atov\\'a, 1990, 2009;  Charv\\'atov\\'a and St\\v{r}e\\v{s}t\\'ik, 2004; Mackey, 2007; Wilson \\emph{et al.}, 2008; Hung, 2007]   suggested that solar variation may be driven by the  planets through  gravitational  spin-orbit coupling mechanisms and tides. These authors have used the inertial motion of the Sun around the center of mass of the solar system (CMSS) as a proxy for describing this phenomenon.  Then, a varying Sun would influence the climate by means of several and complicated mechanisms and feedbacks [Idso and Singer, 2009]. Indeed,  tidal patterns on the Sun well correlate with large solar flare occurrences, and the alignment of Venus, Earth and Jupiter well synchronizes with the 11-year Schwabe solar cycle [Hung, 2007].  In addition, the Earth-Moon system and the Earth's orbital parameters can also be directly modulated by the planetary oscillating gravitational and magnetic fields, and synchronize with their frequencies [Scafetta, 2010]. The Moon can influence the Earth through gravitational tides and orbital oscillations [Keeling and Whorf, 1997, 2000; Munk and Wunsch, 1998;  Munk  and Bills, 2007].  It could be argued that planetary tidal forces are \\emph{weak} and unlikely have any physical outcome. It can also be argued that the tidal forces generated by the terrestrial planets are comparable or even larger than those induced by the massive jovian planets.  However, this is not a valid physical rationale because still little is  known about the solar dynamics and the terrestrial climate. Complex systems are usually characterized by feedback mechanisms that can   amplify the effects of weak periodic forcings also by means of resonance and collective synchronization processes [Kuramoto, 1984; Strogatz, 2009].  Thus, unless the physics of a system is not clearly understood, good empirical correlations at multiple time scales cannot be dismissed just because the microscopic physical mechanisms may be still obscure and need to be investigated.  The above theory implies the existence of direct and/or indirect links between the motion of the planets and the climate oscillations, essentially claiming that the climate is synchronized to the natural oscillations of the solar system, which are driven  by the movements of the planets around the Sun.  If this theory is correct, it can be efficiently used for interpreting climate changes and forecasting climate variability because the motion of the planets can be rigorously calculated.  In this paper we investigate this theory by testing a synchronization hypothesis, that is, whether the planetary motion and the climate present a common set of frequencies. Further, we compare the statistical performance of a phenomenological planetary model for interpreting the climate oscillations with that of a typical major general circulation model adopted by the IPCC. Our findings show that a planetary-based climate model would largely outperform the traditional one in reconstructing the oscillations observed in the climate records. ",
        "conclusions": "On secular, millenarian and larger time scales astronomical oscillations and solar changes drive climate variations.  Shaviv's theory [2003] can explain the large 145 Myr climate oscillations during the last 600 million years.  Milankovic's theory [1941] can explain the multi-millennial climate oscillations observed during the last 1000 kyr. Climate oscillations with periods of 2500, 1500, and 1000 years during the last 10,000 year (the Holocene) are correlated to equivalent  solar cycles that caused the Minoan, Roman, Medieval and Modern warm periods [Bond \\emph{et al.}, 2001; Kerr, 2001]. Finally, several other authors found that multisecular  solar oscillations  caused bi-secular little ice ages (for example: the Sp\\\"orer, Maunder, Dalton minima) during the last 1000 years [for example: Eddy, 1976; Eichler \\emph{et al.}, 2009; Scafetta and  West, 2007;  Scafetta, 2009, 2010].   Herein, we have found empirical evidences that the climate oscillations within the secular scale are very likely driven by astronomical cycles, too. Cycles with periods of 10-11, 12, 15, 20-22, 30 and 60 years are present in all major surface temperature records since 1850, and can be easily linked to the orbits of Jupiter and Saturn. The 11 and 22-year cycles are the well-known Schwabe and Hale solar cycles.  Other faster cycles with periods between 5 and 10 years are in common between the temperature records and the astronomical cycles. Long-term lunar cycles induce a 9.1-year cycle in the temperature records and probably other cycles, including an 18.6-year cycle in some regions [McKinnell and Crawford, 2007]. A quasi-60 year cycle has been found in numerous multi-secular climatic records, and it is even present in the traditional Chinese, Tibetan and Tamil calendars, which are arranged in major 60-year cycles.    The physical mechanisms that would explain this result are still unknown. Perhaps the four jovian planets modulate solar activity via gravitational and magnetic forces that cause tidal and angular momentum stresses on the Sun and its heliosphere. Then, a varying Sun modulates climate, which amplifies the effects of the solar input through several feedback mechanisms. This phenomenon is mostly regulated by Jupiter and Saturn, plus some important contribution from Neptune and Uranus, which modulate a bi-secular cycle with their 172 year synodic period. This interpretation is supported by the fact that the 11-year solar cycles and the solar flare occurrence appear synchronized to the tides generated on the Sun by Venus, Earth and Jupiter [Hung, 2007]. Moreover, a 60-year cycle and other planetary cycles have been found in millennial solar records [Ogurtsov \\emph{et al.}, 2002] and in the number of middle latitude auroras [Komitov, 2009].    Alternatively, the planets are directly influencing the Earth's climate by modulating the orbital parameters of the Earth-Moon system and of the Earth. Orbital parameters can modulate the Earth's angular momentum via gravitational tides and magnetic forces. Then, these orbital oscillations are amplified by the climate system through synchronization of its natural oscillators. This interpretation is supported by the fact that the temperature records contain a clear 9.1-year cycle, which is associated to some long-term lunar tidal cycles. However, the climatic influence of the Moon may be more subtle because several planetary cycles are also found in the Earth-Moon system.     The astronomical forcings may be modulating the length of the day (LOD). LOD presents a 60-year cycle that anticipates the 60-year temperature cycle [Klyashtorin 2001; Klyashtorin and Lyubushin,  2007, 2009; Mazzarella, 2007, 2008; Sidorenkov and Wilson, 2009]. A LOD change  can drive the ocean oscillations by exerting some pressure on the ocean floor and by modifying the Coriolis' forces. In particular, the large ocean oscillations such as the AMO and PDO oscillations are likely driven by astronomical oscillations.    The results herein found show   that the climate oscillations are driven by multiple astronomical mechanisms. Indeed, the planets with their movement cause the entire solar system to vibrate with a set of frequencies that are closely related to the orbital periods of the planets. The wobbling of the Sun around the center of mass of the solar system is just the clearest manifestation of these solar system vibrations and has been used herein just as a proxy for studying those vibrations. The Sun, the Earth-Moon system and the Earth feel these oscillations, and it is reasonable that the internal physical processes of the Earth and the Sun synchronize to them.   It is evident that we can still infer, by means of a detailed data analysis, that the solar system likely induces the climate oscillations, although the actual mechanisms that explain the observed climate oscillations are still unknown. If the true climate mechanisms were already known and well understood, the general circulation climate models would properly reproduce the climate oscillations. However, we found that this is not the case. For example, we showed that the GISS ModelE fails to reproduce the climate oscillations at multiple time scales, including the large 60-year cycle. This failure is common to all climate models adopted by the IPCC [2007] as it is evident in their figures 9.5 and SPM.5 that show the multi-model global average simulation of surface warming. This failure indicates that the models on which the IPCC's claims are based are still incomplete and possibly flawed.      The existence of a 60-year natural cycle in the climate system, which is clearly proven in multiple studies and herein in Figures 2, 6, 10 and 12, indicates that the AGWT promoted by the IPCC [2007], which claims that 100\\% of the global warming observed since 1970 is anthropogenic, is erroneous.   In fact, since 1970 a global warming of about 0.5 $^oC$ has been observed. However, from 1970 to 2000 the 60-year natural cycle was in his warming phase and has contributed no less than 0.3 $^oC$ of the observed 0.5 $^oC$ warming, as Figure 10B shows. Thus, at least 60\\% of the observed warming since 1970 has been naturally induced.   This leaves less than 40\\% of the observed warming to human emissions. Consequently, the current climate models, by failing to simulate the observed quasi-60 year temperature cycle, have significantly overestimated the climate sensitivity to anthropogenic GHG emissions by likely a factor of three. Moreover, the upward trend observed in the temperature data since 1900 may be partially due to land change use, uncorrected urban heat island effects [McKitrick and Michaels, 2007; McKitrick, 2010] and to the bi-secular and millennial solar cycles that reached their maxima during the last decades [Bond \\emph{et al.}, 2001; Kerr, 2001; Eichler \\emph{et al.}, 2009; Scafetta, 2010].   Solomon \\emph{et al.} [2010] recently acknowledged that stratospheric water vapor, not just anthropogenic GHGs, is a very important climate driver of the decadal global surface climate change.  Solomon \\emph{et al.} estimated that stratospheric water vapor has largely contributed both to the warming observed from 1980-2000 (by 30\\%) and to the slight cooling observed after 2000 (by 25\\%). This study reinforces that climate change is more complex than just a response to added $CO_2$ and a few other anthropogenic GHGs. The causes of stratospheric water vapor variation are not understood yet. Perhaps, stratospheric water vapor is driven by UV solar irradiance variations through ozone modulation, and works as a climate feedback to solar variation [Stuber \\emph{ et al.}, 2001]. Thus, Solomon's finding would partially support the findings of this paper and those of Scafetta and West [2005, 2007] and Scafetta [2009]. The latter studies found a significant natural and solar contribution to the warming from 1970-2000 and to the cooling afterward.   A detailed reconstruction of the climate oscillations suggests that a   model based on celestial oscillations, as shown in Figure 12, would largely outperform current general circulation climate models, such as the GISS ModelE, in reconstructing the climate oscillations. The planetary model would also be  more accurate in forecasting climate changes during the next few decades. Over this time, the global surface temperature will likely remain approximately steady, or actually cool.    In conclusion, data analysis indicates that current general circulation climate models are missing fundamental mechanisms that have their physical origin and ultimate justification in astronomical phenomena, and in interplanetary and solar-planetary interaction physics."
    },
    "1005/1005.2425_arXiv.txt": {
        "abstract": "{ We study how the universe reheats following an era of chaotic Dirac--Born--Infeld inflation, and compare the rate of particle production with that in models based on canonical kinetic terms. Particle production occurs through non-perturbative resonances whose structure is modified by the nonlinearities of the Dirac--Born--Infeld action. We investigate these modifications and show that the reheating process may be efficient. We estimate the initial temperature of the subsequent hot, radiation-dominated phase.}  \\preprint{arXiv:1005.2425} ",
        "introduction": "\\label{sec:introduction}  Measurements of the microwave background temperature made by the WMAP satellite over the last decade \\cite{Komatsu:2010fb} are consistent with the idea that structure in the universe originated as small fluctuations during a primordial era of accelerated expansion, or `inflation'. Inflation leaves the universe in a cold vacuum state, devoid of radiation and matter. If our own universe experienced an inflationary era in its past, we must conclude that some process repopulated its sterile vacuum with the abundant quanta we see today. Much less is known about this `reheating' process than is understood about the observationally accessible density fluctuation.  Properties of inflationary fluctuations largely decouple from the details of whatever physics drives the accelerating epoch, leaving only weak traces.% \\footnote{These traces are presently the subject of intense theoretical and observational effort; see, e.g. Ref.~\\cite{Chen:2010xk}.} This freedom has encouraged many proposals for the underlying microphysics, nearly all of which are equally unconstrained by observation. In an especially interesting proposal, due to Silverstein \\& Tong \\cite{Silverstein:2003hf} and later elaborated by Alishahiha, Silverstein \\& Tong \\cite{Alishahiha:2004eh}, inflation occurs while a $\\D{3}$-brane moves in a warped background spacetime.% \\footnote{D-branes are extended objects in spacetime. For a review of their properties see Johnson \\cite{Johnson:2003gi}; inflation in such models has been reviewed by McAllister \\& Silverstein \\cite{McAllister:2007bg}. Many detailed properties were discussed by Chen~\\cite{Chen:2005ad}.} As we shall discuss in \\S\\ref{sec:review}, the motion of the brane is controlled by an action of Dirac--Born--Infeld (``DBI'') type and the resulting model is known as DBI inflation. In this paper we study the reheating era which must follow such an inflationary phase. Reheating in related brane-world models has been studied in Refs.~\\cite{Cline:2002it, Brodie:2003qv, Dvali:2003ar, Frey:2005jk, Chialva:2005zy}. In contrast to the primordial perturbations, reheating depends strongly on many model-specific details~\\cite{Kofman:1996mv,Kaiser:1995fb}. It is a subtle function of nonlinear physics, probing information complementary to the linear physics of the density fluctuation. Nonlinear properties vary widely between competing inflationary models, and there are important differences between DBI models and those with canonical kinetic terms.  Many authors have contributed to the theory of particle production, which is now well-developed \\cite{Traschen:1990sw,Kofman:1994rk,Shtanov:1994ce,Kofman:1997yn}.% \\footnote{The theory of preheating has been reviewed by Allahverdi {\\etal} \\cite{Allahverdi:2010xz}, Kofman {\\etal} \\cite{Kofman:2008zz} and Bassett {\\etal} \\cite{Bassett:2005xm}. See also Ref.~\\cite{Kohri:2009ac}.} In \\emph{perturbative reheating}, the decay rate of inflaton particles into other species of matter is calculated directly from an S-matrix. These calculations assume any relics of previous decays to be sufficiently diffuse that particle production takes place into an effectively pristine vacuum. It was later understood that a Bose enhancement of the decay rate may occur for integer-spin species if the products of previous decays accumulate, leading to a resonant phase of out-of-equilibrium production known as \\emph{preheating}. A different effect, so-called \\emph{tachyonic preheating} \\cite{Felder:2000hj,Felder:2001kt, Dufaux:2006ee}, is associated with the end of inflation in certain models. Unlike conventional preheating this does not merely convert one species of particle into another, but (like any tachyonic instability) reflects a preference to convert one vacuum state to another by the rapid accumulation and condensation of particles.  These phenomena may occur in any effective field theory of the post-inflationary universe, whatever its microphysical origin. The general theory has subsequently been applied to concrete models, including some versions of brane inflation, in which it is possible to give a more refined interpretation \\cite{Kachru:2003sx,Barnaby:2004gg}. In these models, inflation ends when the moving $\\D{3}$-brane becomes close to an antibrane (a $\\Dbar{3}$-brane), allowing open strings stretching between the $\\D{3}$--$\\Dbar{3}$ pair to become excited. The open string states include a tachyon which induces fragmentation of the $\\D{3}$ brane into $\\D{0}$ branes, described on the $\\D{3}$ worldvolume by a phase of tachyonic preheating. The $\\D{0}$ fragments decay into closed string states, which are enumerated by the Kaluza--Klein (``KK'') modes of supergravity fields in the warped background. These KK modes subsequently decay into the species of matter and radiation which populate our universe \\cite{Barnaby:2004gg}.  In this paper we study a version of the DBI model in which the end of inflation is less dramatic. We work in a model where the $\\D{3}$ comes to rest near an extremity of its warped background, referred to as the ``tip of the throat.'' As it settles on its resting-place it executes coherent oscillations, inducing growing fluctuations which fold and wrinkle its surface. This scenario was suggested by Silverstein \\& Tong \\cite{Silverstein:2003hf}. Our interest is in the \\emph{qualitative} differences which may emerge in concrete models. Should we expect the reheating temperature to increase or decrease? Is parametric resonance more efficient, perhaps due to the extra non-linearities of the DBI action, or do these same non-linearities conspire to shut off the resonant channels? For this reason we focus on differences with the case of  canonical kinetic terms and largely ignore similarities. For example, the use of perturbation theory in either case will break down at the onset of back-reaction, when the population of matter species grows to the point where it cannot be ignored. This affects the canonical and DBI cases equally, and we restrict our comparison to times before back-reaction becomes significant.  This paper is organized as follows. In \\S\\ref{sec:review} we review the dynamics of inflation in DBI models and discuss a toy model in which inflation ends as the brane settles to its final position in the warped background. This model is comparable to chaotic inflation in a theory with canonical kinetic terms. In \\S\\ref{sec:res} we discuss parametric resonance, making a comparison with canonical chaotic inflation. In \\S\\ref{sec:pr} we study the stage of reheating and estimate the reheating temperature. We conclude in \\S\\ref{sec:conclude}.  Throughout this paper, we work in units where $\\hbar = c = 1$ and set the reduced Planck mass to unity, $\\Mp \\equiv (8 \\pi G)^{-1/2} = 1$. The metric signature is taken to be $(-, +, +, +)$. ",
        "conclusions": "\\label{sec:conclude}  We have studied mechanisms by which the universe can be repopulated with matter species following an era of D-brane inflation. We work in the limit where the motion of the brane is at most perturbatively relativistic, which may ultimately be required by observation if $\\fnl$ measured on equilateral configurations of the bispectrum is nonnegative.  Our analysis is valid in a regime where the amplitude of oscillations is significantly smaller than whatever infrared scale caps the throat in which the D-brane moves. Indeed, because we take the D-brane to be at most perturbatively relativistic, inflation is of the slow-roll type. We assume inflation ends as the D-brane gently coasts to its resting position, dissipating its remaining energy by executing small amplitude oscillations. The amplitude and phase of these oscillations are mildly perturbated in comparison with the oscillations of canonically normalized scalar reheating the universe after a phase of four-dimensional inflation.  We encounter some novel effects which are not present in a theory with canonical kinetic terms. In canonical inflation there is no resonant production of inflaton particles during the reheating era. By contrast, in the perturbatively-relativistic limit of DBI inflation, non-linearities arising from the Dirac--Born--Infeld action cause a weak resonance which reaches an $\\Or(1)$ effect only after $\\sim 1/\\epsilon$ oscillations. Excessive production of inflaton particles is typically problematic, but since preheating typically concludes after $\\lesssim 10$ oscillations this weak resonance is unlikely to be fatal for the DBI model. In the ultra-relativistic limit where $\\gamma \\gg 1$, the potential becomes unimportant and presumably resonance will be strongly suppressed.  Another novel effect is a `DBI friction,' which appears in the equation of motion for the inflaton. This slightly damps the preheating effect. Beyond perturbation theory, it is possible this effect represents the decreasing significance of the potential at large $\\gamma$.  For a given set of measured masses and couplings which characterize the properties of the inflaton and the matter field into which it decays, we find that final reheating temperatures (assuming they are achieved by perturbative decays) are typically smaller than in the equivalent theory with canonical kinetic terms. This may be beneficial in concrete models, where lower reheating temperatures allow problems associated with overclosure of the universe by moduli or gravitinos to be ameliorated. The most stringent constraint on DBI models comes from observations of the microwave background bispectrum, which presently require $\\gamma \\lesssim 18$. If our results are representative, and the reheating temperature falls as $\\gamma$ increases, a constraint on $\\gamma$ may also emerge from demanding that the nucleosynthesis era is uninterrupted, which requires $T_{\\mathrm{RH}} \\gtrsim \\mbox{MeV}$."
    },
    "1005/1005.2563_arXiv.txt": {
        "abstract": "{Within the framework of the HERM33ES key project, we are studying the star forming interstellar medium in the nearby, metal-poor spiral galaxy M33, exploiting the high resolution and sensitivity of \\herschel.  } {We use PACS and SPIRE maps at 100, 160, 250, 350, and 500\\,$\\mu$m wavelength, to study the variation of the spectral energy distributions (SEDs) with galacto-centric distance. } {Detailed SED modeling is performed using azimuthally averaged fluxes in elliptical rings of 2\\,kpc width, out to 8\\,kpc galacto-centric distance. Simple isothermal and two-component grey body models, with fixed dust emissivity index, are fitted to the SEDs between 24\\,$\\mu$m and 500$\\,\\mu$m using also MIPS/\\spitzer\\, data, to derive first estimates of the dust physical conditions. } {The far-infrared and submillimeter maps reveal the branched, knotted spiral structure of M33. An underlying diffuse disk is seen in all SPIRE maps (250--500\\,$\\mu$m). Two component fits to the SEDs agree better than isothermal models with the observed, total and radially averaged flux densities.  The two component model, with $\\beta$ fixed at 1.5, best fits the global and the radial SEDs. The cold dust component clearly dominates; the relative mass of the warm component is less than 0.3\\% for all the fits. The temperature of the warm component is not well constrained and is found to be about 60\\,K$\\pm$10\\,K. The temperature of the cold component drops significantly from $\\sim24$\\,K in the inner 2\\,kpc radius to 13\\,K beyond 6\\,kpc radial distance, for the best fitting model. The gas-to-dust ratio for $\\beta=1.5$, averaged over the galaxy, is higher than the solar value by a factor of 1.5 and is roughly in agreement with the subsolar metallicity of M33. } {} ",
        "introduction": "%  In the local universe, most of the observable matter is contained in stellar objects that shape the morphology and dynamics of their ``parent\" galaxy. In view of the dominance of stellar mass, a better understanding of star formation and its consequences is mandatory. There exists a large number of high spatial resolution studies related to individual star forming regions of the Milky Way, as well as of low linear resolution studies of external galaxies. For a comprehensive view onto the physical and chemical processes driving star formation and galactic evolution it is, however, essential to combine local conditions affecting individual star formation with properties only becoming apparent on global scales.    At a distance of 840\\,kpc \\citep{freedman1991}, M33 is the only nearby, gas rich disk galaxy that allows a coherent survey at high spatial resolution. It does not suffer from any distance ambiguity, as studies of the Milky Way do, and it is not as inclined as the Andromeda galaxy. M33 is a regular, relatively unperturbed disk galaxy, as opposed to the nearer Magellanic Clouds, which are highly disturbed irregular dwarf galaxies.  M33 is among the best studied galaxies; it has been observed extensively at radio, millimeter, far-infrared (FIR), optical, and X-ray wavelengths, ensuring a readily accessible multi-wavelength database. These data trace the various phases of the interstellar medium (ISM), the hot and diffuse, the warm and atomic, as well as the cold, dense, star forming phases, in addition to the stellar component.  However, submillimeter and far-infrared data at high angular and spectral resolutions have been missing so far.   In the framework of the open time key project ``\\herschel\\, M33 extended survey ({\\tt HERM33ES})'', we use all three instruments onboard the ESA \\herschel\\, Space Observatory \\citep{pilbratt2010} to study the dusty and gaseous ISM in M33. One focus of {\\tt HERM33ES} is on maps of the FIR continuum observed with PACS \\citep{poglitsch2010} and SPIRE \\citep{griffin2010}, covering the entire galaxy.  A second focus lies on observing diagnostic FIR and submillimeter cooling lines \\CII, \\OI, \\NII, and H$_2$O, toward a $2'\\times40'$ strip along the major axis with PACS and HIFI \\citep{degraauw2010}.    \\begin{figure}[t] \\centering \\includegraphics[width=7cm]{14613fg1.eps} \\caption{A composite 500\\,$\\mu$m (red) and 160$\\,\\mu$m (blue) map of M33. The most extended emission is traced by the longest wavelength map revealing the presence of the cold dust in the outskirts of the galaxy.  The shorter wavelengths mostly trace the branched spiral structure as well as distinct warm \\HII\\ regions and star-forming complexes (such as the ones labeled).} \\label{fig-maps} \\end{figure}      In this first {\\tt HERM33ES} paper, we use continuum maps covering the full extent of M33, at 100, 160, 250, 350, and $500\\,\\mu$m. These data are an improvement over previous data sets of M33, obtained with ISO and \\spitzer\\, \\citep{hippelein2003,hinz2004,tabatabaei2007-466}, in terms of wavelength coverage and angular resolution.  The total bolometric luminosity of normal galaxies is only about a factor of 2 larger than the total IR continuum emission \\citep{hauser-dwek2001}, which in turn accounts for more than $\\sim98$\\% of the emission of the ISM (dust$+$gas) \\citep[e.g.][]{malhotra2001,dale2001}. Massive star formation heats the dust mainly via its far-ultraviolet (FUV) photons and the absorbed energy is then reradiated in the IR. FIR continuum fluxes are therefore often used as a measure of the interstellar radiation field (ISRF) \\citep[e.g.][]{kramer2005} and the star formation rate (SFR) \\citep[e.g.][]{schuster2007}.  However, a number of authors have suggested that half of the FIR emission or more is due to dust heated by a diffuse ISRF, and not directly linked to massive star formation \\citep{israel1996, verley2009}.     Another disputed topic is the evidence for a massive, cold dust component in galaxies. The SCUBA Local Universe Galaxy Survey \\citep{dunne2001} identified a cold dust component at an average temperature of 21\\,K. A number of studies of the millimeter continuum emission of galaxies found indications for even lower temperatures \\citep{misiriotis2006,weiss2008,liu2010}. In order to estimate the amount of dust at temperatures below about 20\\,K, and to improve our understanding of the physical conditions of the big grains, well calibrated observations longward of $\\sim150\\,\\mu$m wavelength are needed. ",
        "conclusions": ""
    },
    "1005/1005.5343_arXiv.txt": {
        "abstract": "The multiplicity distributions produced by the variation of time-dependent gravitational fields in a conformally flat background geometry belong to the same class of infinitely divisible distributions found, for fixed centre of mass energies and symmetric (pseudo)rapidity intervals, in charged multiplicities produced in $pp$, $p\\overline{p}$ and in heavy ion collisions. Apparently unrelated multiplicity distributions are classified in terms of the (positive) discrete representations of the $SU(1,1)$ group. The gravitational analogy suggest a global high-energy asymptote for the distributions measured in $pp$ and $p\\overline{p}$ collisions. Second-order cross correlations between positively and negatively charged distributions represent a relevant diagnostic for a closer scrutiny of the multiparticle final state. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.3279_arXiv.txt": {
        "abstract": "{We present far-infrared and submillimetre spectra of three carbon-rich evolved objects, AFGL 2688, AFGL 618 and NGC 7027. The spectra were obtained with the SPIRE Fourier transform spectrometer on board the {\\it Herschel} Space Observatory, and cover wavelengths from 195-670$\\mu$m, a region of the electromagnetic spectrum hitherto difficult to study in detail.  The far infrared spectra of these objects are rich and complex, and we measure over 150 lines in each object.  Lines due to 18 different species are detected.  We determine physical conditions from observations of the rotational lines of several molecules, and present initial large velocity gradient models for AFGL 618.  We detect water in AFGL~2688 for the first time, and confirm its presence in AFGL~618 in both ortho and para forms.  In addition, we report the detection of the J=1-0 line of CH$^+$ in NGC~7027.} {}{}{}{} ",
        "introduction": "Low- to intermediate-mass stars ($<$8M$_{\\odot}$) shed much of their mass during the final stages of their evolution, in the form of a slow molecular wind.  The mass loss causes the star to leave the Asymptotic Giant Branch (AGB), moving to the left in the Hertzsprung-Russell diagram.  The star becomes hotter, warming and eventually ionising its previously ejected outer layers, which become visible as a planetary nebula (PN).  The ejecta from such stars are a major contributor to galactic chemical evolution (e.g. Matsuura et al. 2009).  The evolution from the AGB to the PN stage is rapid, and relatively few objects in this intermediate stage are known (326 are listed by Szczerba et al. (2007), compared to around 3000 known or suspected PNe (Frew \\& Parker 2010)).  AFGL~618 and AFGL~2688 (the Egg Nebula) are two of the best-known objects in this transition stage.  AFGL~2688 is illuminated by a central star with spectral type F5, and Hubble Space Telescope images reveal a large number of round arcs crossed by `searchlight beams' (Sahai et al. 1998).  Proper motion measurements give a distance of 420pc and a dynamical age for the ejected material of 350 years (Ueta et al. 2006).  AFGL~618 entered the protoplanetary nebula (PPN) phase 100-200 years ago (Kwok \\&  Bignell 1984, Bujarrabal et al. 1988), and is more evolved than AFGL~2688.  It contains a B0 central star surrounded by a compact H~{\\sc ii} region (Wynn-Williams 1977, Kwok \\&  Bignell 1984).  The variability of free-free emission from this object over the last 30 years implies rapid evolution (Kwok \\&  Feldman 1981, S\\'anchez Contreras et al. 2002).  NGC~7027 is a very young planetary nebula (Masson 1989).  It has been extensively studied due to its high surface brightness at all wavelengths, and observations with the Infrared Space Observatory (ISO) revealed a far-infrared spectrum rich in atomic and molecular lines (Liu et al. 1996).  Its ionised inner regions are surrounded by and partly obscured by a PDR and massive molecular envelope.  Previous far infrared and sub-mm observations of these evolving post-AGB objects, although limited in their spatial resolution and spectral coverage, have revealed rich spectra containing numerous molecular and dust features (e.g. Cox et al., 1996, Liu et al. 1996, Herpin et al. 2002, Pardo et al. 2004, 2005, 2007a, 2007b).  The {\\it Herschel} Space Observatory (Pilbratt et al. 2010) now significantly extends our observational capabilities.  The SPIRE spectrometer covers wavelengths from 195-670$\\mu$m, a region of the electromagnetic spectrum hitherto largely unexplored.  The SPIRE instrument, its in-orbit performance, and its scientific capabilities are described by Griffin et al. (2010), and the SPIRE astronomical calibration methods and accuracy are outlined by Swinyard et al. (2010).  Here we present observations of AFGL~2688, AFGL~618 and NGC~7027 using the SPIRE spectrometer. ",
        "conclusions": "The unprecedented wavelength coverage and excellent sensitivity of the SPIRE spectrometer have allowed us to see the incredibly rich far-infrared and submm spectra of three carbon-rich evolved stars for the first time.  We have detected 150-200 emission lines in the SPIRE FTS spectra of AFGL~2688, AFGL~618 and NGC~7027, and presented preliminary investigations of the physical conditions in these objects based on the new data.  We find water in both ortho and para forms in all three objects, and observe strong ionised lines in the more evolved NGC 7027, including the J=1-0 line of CH$^{+}$.  This line, never previously detected in an astronomical source, is also detected in {\\it Herschel} spectra of H~{\\sc ii} region PDRs (Naylor et al. 2010).  LVG models of AFGL~618 are in excellent agreement with those based on earlier ISO data.  These observations show that the SPIRE spectrometer on board {\\it Herschel} can provide significant new insights into the temperatures structure and chemical content of the outflows from evolved stars."
    },
    "1005/1005.4459_arXiv.txt": {
        "abstract": "{Classical T Tauri stars (CTTS) differ in their X-ray signatures from older pre-main sequence stars, e.g. weak-lined TTS (WTTS). CTTS show a soft excess and deviations from the low-density coronal limit in the He-like triplets.}{We test whether these features correlate with accretion or the presence of a disk by observing IM~Lup, a disk-bearing object apparently in transition between CTTS and WTTS without obvious accretion.}{We analyse a \\emph{Chandra} grating spectrum and additional \\emph{XMM-Newton} data of IM~Lup and accompanying optical spectra, some of them taken simultaneously to the X-ray observations. We fit the X-ray emission lines and decompose the H$\\alpha$ emission line in different components.}{In X-rays IM~Lup has a bright and hot active corona, where elements of low first-ionisation potential are depleted. The He-like \\ion{Ne}{ix} triplet is in the low-density state, but due to the small number of counts a high-density scenario cannot be excluded on the 90\\% confidence level. In all X-ray properties IM~Lup resembles a  main-sequence star, but it is also compatible with CTTS signatures on the 90\\% confidence level, thus we cannot decide if the soft excess and deviations from the low-density coronal limit in the He-like triplets in CTTS require accretion or only the presence of a disk. IM~Lup is chromospherically active, which explains most of the emission in H$\\alpha$. Despite its low equivalent width, the complexity of the H$\\alpha$ line profile is reminiscent of CTTS. We present an estimate for the mass accretion rate of $10^{-11}M_{\\sun}$~yr$^{-1}$. }{} ",
        "introduction": "\\label{introduction} Our view on the formation of stars and planetary systems and their emergence from molecular clouds has made significant progress over the last decades. Stars form when giant molecular clouds fragment and contract to form proto-stars. Mass accretion onto those stellar cores proceeds via an accretion disk, while the surrounding envelope eventually disperses. The low-mass pre-main sequence stars in this stage come in two types, the classical T~Tauri stars (CTTS) and the weak-lined T~Tauri stars (WTTS). Traditionally, they were distinguished only by their H$\\alpha$ emission line equivalent width, with those stars of EW $> 10$~\\AA{} defined to be CTTS. It turned out that the EW of H$\\alpha$ is a good tracer of the accretion flow from the circum-stellar disk the CTTS still possess. The H$\\alpha$ EW of WTTS is smaller and their line profiles are symmetric; in contrast CTTS exhibit broader emission lines, which are sometimes asymmetric \\citep{2009A&A...504..461F}. Usually the combination of strong H$\\alpha$ emission and line asymmetry is a reliable accretion indicator \\citep{1998ApJ...492..743M,2003ApJ...592..266M}.  WTTS with their low H$\\alpha$ EW are generally expected not to show accretion. It is suggestive to interpret the WTTS as more evolved CTTS, where accretion has stopped already. This does not necessarily imply the absence of a disk, as transitional disks may exist without ongoing accretion for some time \\citep{2006ApJ...645.1283P}. For many TTS the H$\\alpha$ EW is known to be variable, possibly because the accretion switches on and off or --at least-- the accretion rate changes.  Both types of TTS have been known for a long time to be copious X-ray emitters \\citep{1999ARA&A..37..363F}. As a class CTTS stand out from other X-ray sources by their strong soft X-ray excess \\citep{RULup,manuelnh}, but CTTS also exhibit normal coronal activity and stellar flares.  The CTTS TW~Hya, observed with \\emph{Chandra}/HETGS and \\emph{XMM-Newton} \\citep{2002ApJ...567..434K,twhya}, was the first star where unusual line ratios in the He-like triplets was found indicating high densities in the formation region. This phenomenon repeats in most CTTS observed so far (\\object{BP Tau}: \\citet{bptau}; \\object{V4046 Sge}: \\citet{v4046}; \\object{RU Lup}: \\citet{RULup}; \\object{MP Mus}: \\citet{2007A&A...465L...5A}; \\object{Hen 3-600}: \\citet{2007ApJ...671..592H}). There are exceptions to this rule -- in the more massive, eponymous T~Tau itself, although known for its high accretion rate, the \\ion{O}{vii} triplet is consistent with the coronal limit \\citep{ttau}. The same has been found for HAeBe stars, which are in a similar evolutionary state as the CTTS, yet of higher mass (\\object{AB Aur}: \\citet{ABAur}; \\object{HD 163296}: \\citet{HD163296}). The soft excess and the He-like triplet line ratios in CTTS can be naturally linked to mass accretion from proto-planetary disks \\citep{lamzin,acc_model}.   Systems containing typical WTTS such as \\object{TWA 4} \\citep{2004ApJ...605L..49K} and \\object{TWA 5} \\citep{twa5} show no sign of active accretion and no high densities in their X-ray spectra. Well studied examples of young main-sequence (MS) stars such as AU~Mic, Speedy~Mic and AB~Dor \\citep{2004A&A...427..667N} also show triplets indicating low densities similar to older MS objects. Still, 5 of 83 WTTS observed by \\citet{2006ApJ...645.1283P} do show an IR-excess caused by dust within a few AU from the star, detected by \\emph{Spitzer}, thus there is a matter reservoir for accretion.  Observationally, we could not distinguish in the past if the X-ray signatures observed in CTTS in contrast to WTTS and young MS stars are caused by active accretion or the presence of a disk. To solve this question we observed a system with \\emph{Chandra}, that was previously classified as disk-bearing, non-accreting WTTS. Simultaneous to the X-ray observations, the accretion state is observed using optical spectroscopy.   The properties of the our target IM~Lup are summarised in Sect.~\\ref{imlupprop}. We then present the X-ray observations of IM~Lup and the accompanying optical data in Sect.~\\ref{observations}. We show our new results and compare IM~Lup to CTTS, WTTS and MS stars in Sect.~\\ref{results}. In Sect.~\\ref{discussion} we discuss the implications and end with a short summary in Sect.~\\ref{summary}. ",
        "conclusions": "\\label{discussion}  \\subsection{X-ray properties} IM~Lup shows many characteristics of an active star. It has a bright, hot corona and its abundances follow the IFIP pattern. These traits are common to both CTTS and WTTS. In the density diagnostic of the He-like triplets there is no deviation from the coronal limit, however, due to the low signal a high-density state is allowed at the 90\\% confidence level. The signal of the \\ion{O}{vii} triplet is too weak to check for any soft X-ray excess. Thus, we conclude that -- from the X-ray point-of-view -- it is more likely that IM~Lup shares its characteristics with WTTS, despite the presence of its disk. If confirmed, this suggests that the distinctive characteristics of CTTS are really due to accretion: An excess of soft X-ray emission originates in the high-density environment of an accretion spot as simulated by \\citet{lamzin} and \\citet{acc_model}. Our observation is incompatible with any hypothetical scenario where the high-density signatures of the He-like triplets are attributed to the accretion disk itself because they are not observed (on the 1$\\sigma$ confidence level) in IM Lup, which has a disk.  The only peculiarity in comparison to the WTTS samples of COUP and XEST is the high bolometric luminosity of IM~Lup. It belongs to the brightest objects of its kind. This may be due to an overestimation of the distance. We use $190\\pm27$~pc from \\citet{1998MNRAS.301L..39W} based on the \\emph{HIPPARCOS} parallax, but other methods lead to lower values of $140\\pm20$~pc \\citep{1993AJ....105..571H}. Both values are compatible within the $2\\sigma$ errors, but if the lower value was the true one, X-ray emission measure and bolometric luminosity would be only half of the values given above. This does not change any of the other conclusions of our analysis, but IM~Lup would fit in better with the bulk of the COUP and XEST objects in this case. Still, IM~Lup is an active star with a bright and hot corona and $\\log (L_X/L_{bol})=-3.1$, similar to saturated MS stars and also typical for CTTS \\citep{2009A&ARv..17..309G}. %   \\subsection{Mass accretion?} \\label{massacretion} IM~Lup has a high level of chromospheric activity, which causes the emission observed in H$\\alpha$. Other M dwarfs with active chromosphere also show H$\\alpha$ emission, typically with an equivalent width of a few~\\AA{} \\citep{1981ApJS...46..159W, 1990ApJS...74..891R}, very similar to IM~Lup. Their Ca~H and K lines often show a very narrow emission peak, too. The H$\\alpha$ line profiles vary, but in  most active stars there is a single, narrow emission peak, sometimes with a reversal due to self-absorption in the centre. This is in line with the central peak in the IM~Lup data. In the samples of \\citet{1981ApJS...46..159W} and \\citet{1990ApJS...74..891R} there is not a single case of H$\\alpha$ in emission, where simultaneously broad absorption is present. The same is true for WTTS, which have narrow emission lines of symmetric shape \\citep{2009A&A...504..461F}.  In contrast, red-shifted absorption dips and red- and blue-shifted emission components appear in accretion models in most viewing geometries \\citep{2001ApJ...550..944M}. Thus, in IM~Lup the H$\\alpha$ line profile must be caused by additional material which is not present in MS stars: Candidates are a circumstellar envelope, the disk or accretion funnels. We do not detect veiling of the photospheric emission lines, thus any mass accretion must be very low, if present at all. Still, the narrow emission components might originate in  a weak accretion funnel. They are shifted by $\\pm100$~km~s$^{-1}$, because we see accretion towards both, the observers side of the star and the back side. The line shift is smaller than usually seen in emission lines from accretion funnels of CTTS, but given the uncertain stellar parameters of IM~Lup this could be due to a combination of low mass, i.e. smaller gravitational potential, and viewing geometry, where the accretion funnel is not parallel to the line-of-sight. %  In some stars the H$\\alpha$ line profile is also influenced by strong flares but this cannot be the case in IM~Lup, because no flare shows up in the X-ray lightcurve.  The study of \\citet{2003ApJ...582.1109W} correlated the width of the H$\\alpha$ line at 10~\\% of the maximum with accretion. They suggest a full width at 10~\\% of the maximum $>270$~km~s$^{-1}$ as borderline between accreting and non-accreting sources, a slightly lower limit of $>200$~km~s$^{-1}$ was presented by \\citet{2003ApJ...592..282J}. The line width of IM~Lup falls just between those two limits; it would be classified as non-accreting by the one and as accreting by the other boundary. \\citet{2004A&A...424..603N} showed that a quantitative relation between line width and mass accretion rate exists, which is valid from CTTS to sub-stellar objects. This relation yields an accretion rate of $10^{-11} M_{\\sun}\\;\\mathrm{yr}^{-1}$, but IM~Lup is chromospherically very active, which explains a large fraction of the H$\\alpha$ flux. Additionally, not all relations employed to convert the H$\\alpha$ flux to a mass accretion rate agree, if accretion is present at all, therefore we treat this number as a rough estimate only.  In summary, the H$\\alpha$ EW classifies IM~Lup as a WTTS, but the H$\\alpha$ line profile is complex as in CTTS. This suggests that mass accretion occurs at a very small rate. The energy released from accretion is too low to influence the X-ray properties significantly. For a mass accretion rate of $10^{-11} M_{\\sun}\\;\\mathrm{yr}^{-1}$, at most 0.5\\% of the X-rays can be due to accretion, assuming that 10\\% of the accretion energy is released in X-rays \\citep{acc_model}. For the interpretation of the X-ray data we can thus safely treat IM~Lup as a non-accreting source.  \\subsection{Varying absorption?} The HARPS H$\\alpha$ line profiles are presented in Fig.~\\ref{ha_harps}. While the fitting with Gaussians (table~\\ref{tabha}) yields an absorption component reaching over the entire line profile, for non-Gaussian components absorption is only required on the red side where the observed line profile falls below the continuum. The situation is different in the ANU data, which was obtained simultaneous to the \\emph{Chandra} observation. No absorption is visible here on the first day and wide absorbtion wings are present on both sides of the line on the second day. We postulate that matter from the disk on its way onto the star causes the H$\\alpha$ absorption. With IM~Lup's mass and radius (Sect.~\\ref{imlupprop}) we calculate the Keplerian rotation of the disk. At the co-rotation radius of $6\\;R_{*}$ it matches the stellar rotation period from Sect.~\\ref{opticalspectra}; in this region the inner disk should be truncated. The absorption component is strong at 100-400~km~s$^{-1}$, which corresponds to the free-fall velocity from the co-rotation radius. Thus the H$\\alpha$ absorbing gas may be found in the gap between star and disk; however, this scenario can only explain the red-shifted part of the absorption, but it strengthens the argument in Sect.~\\ref{massacretion} that the H$\\alpha$ line width at 10\\% maximum could be interpreted as accretion. We speculate, that the blue-shifted absorption in the ANU data is caused by an outflow. CTTS commonly drive winds or even jets, where the mass loss is of the order of 10\\% of the accretion rate \\citep{1990ApJ...354..687C,2008ApJ...689.1112C}. The blue-shifted absorption appears only on the second day of the ANU observation, when the absoprtion on the red side, which we interpret as accretion, is stronger.  We divided the \\emph{Chandra} observation in two parts and fitted the $n_H$ indenpedently to test if the absorption in the X-ray data changes with time. We found no change of the fitted value. Also, the hardness ratio of the X-ray spectra does not change during the observations. This does not rule out variability in $n_H$, but is sets an upper limit. We estimate that an additional absorbing column of $>10^{21}$~cm$^{-2}$ would have been detected. Does the accretion scenario explain why strong H$\\alpha$ absorption coincides with weak X-ray emission? Given the weak limit on variability in $n_H$ and the fact that we have only two datapoints of simultaneous X-ray and optical spectra, the relation between H$\\alpha$ absorption and X-ray luminosity is --at best-- suggestive. Still it is intriguing to interpret the lower X-ray luminosity and the H$\\alpha$ line profiles as absoption by an accretion funnel. This should cause a hardening of the X-ray absorption. This is not observed, but the limit on extra absorption is weak due to the low count rates.   If taken seriously, the relation between X-ray luminosity and H$\\alpha$ absorption fits to the fact that CTTS are on average less luminous in X-rays than WTTS \\citep[e.g.][]{2005ApJS..160..401P,2007A&A...468..443T}.  \\citet{2009ApJ...699L..35D} suggested that it is not the accreting nature of the CTTS that causes their low luminosity; rather, it is the low luminosity that allows the disk to reach closer to the star, i.e., weaker emitters accrete and become CTTS.  \\subsection{IM Lup in context} \\citet{2009A&A...501.1013S} observed variability in optical spectra of \\object{T Cha}. Like IM~Lup T~Cha is also a transition object from CTTS to WTTS without veiling of photospheric lines, i.e. with little or no mass accretion. \\citet{2009A&A...501.1013S} find the H$\\alpha$ profiles to change from emission lines to absorption components over time, simultaneously they observe changes in the reddening and develop a scenario based on inhomogeneous circumstellar extinction similar to our ideas for IM~Lup. They argue this to be caused by the inner parts of the disk, where grain growth takes place and possibly planets form.  Only a small number of CTTS have been observed with X-ray grating spectroscopy, most of them show low $f/i$ ratios in the He-like triplets, but \\object{T Tau} does not. Does IM~Lup confirm  a CTTS population with low densities like T~Tau or does it behave as WTTS? Although the complex H$\\alpha$ line profile points to mass accretion, IM~Lup differs from T~Tau in other characteristics. The main component of the multiple system T~Tau is much more massive, close to the boundary to the HerbigAe/Be stars and it accretes at much higher rates. Whereas IM~Lup's X-ray emission must be coronal, T~Tau could power a significant part of its luminosity from accretion.  We therefore expect IM~Lup to resemble the WTTS and young MS stars in its X-ray properties. Again, only very few WTTS are observed with X-ray gratings, but they all look similar to MS stars. Despite the small number, it is credible that their emission is driven exclusively by a solar-type corona and we find no evidence that the transition object IM~Lup, which still posseses a disk, is different. The transition from a CTTS to a WTTS in X-rays seems to coincide with the decline of accretion as expected from accretion shock models for CTTS. However, the \\ion{Ne}{ix} triplet is the strongest discriminator and the original driver of the reported observation, but at the 90\\% confidence level the CTTS-like high-density scenario cannot be excluded for IM~Lup."
    },
    "1005/1005.4729_arXiv.txt": {
        "abstract": "We analysed thermonuclear (type-I) X-ray bursts observed from the low-mass X-ray binary 4U~1728$-$34 by {\\it RXTE}, {\\it Chandra}\\/ and {\\it INTEGRAL}. We compared the variation in burst energy and recurrence times as a function of accretion rate with the predictions of a numerical ignition model including a treatment of the heating and cooling in the crust. We found that the measured burst ignition column depths are significantly below the theoretically predicted values, regardless of the assumed thermal structure of the neutron star interior. While it is possible that the accretion rate measured by {\\it Chandra}\\/ is underestimated, due to additional persistent spectral components outside the sensitivity band, the required correction factor is typically 3.6 and as high as 6, which is implausible. Furthermore, such underestimation is even more unlikely for {\\it RXTE}\\/ and {\\it INTEGRAL}, which have much broader bandpasses. Possible explanations for the observed discrepancy include shear-triggered mixing of the accreted helium to larger column depths, resulting in earlier ignition, or the fractional covering  of the accreted fuel on the neutron star surface. ",
        "introduction": "\\label{intro}    Thermonuclear (type I) bursts  are triggered by unstable nuclear burning of the material accreted onto the neutron star (NS) surface in low-mass X-ray binary (LMXB) systems \\citep[e.g.,][]{1995xrbi.nasa..175L,2006csxs.book..113S}. The basic theory of type I bursts was outlined shortly after their detection \\citep[e.g.,][]{1976Natur.263..101W,1977Natur.270..310J,1977xbco.conf..127M,1978ApJ...220..291L}. According to these models, the accreted material, which usually consists of hydrogen and helium, accumulates as  a thin layer on the NS surface (typically on time-scales of hours to days), and when the pressure and temperature at its base reach critical values, the fuel will ignite and burn unstably until exhausted.     Further observations \\citep[see][for a review]{2008ApJS..179..360G} and subsequent modeling \\citep{1981ApJ...247..267F,1987ApJ...319..902F,1987ApJ...323L..55F,1998mfns.conf..419B,2003ApJ...599..419N,2004ApJS..151...75W,2006ApJ...652..584C} showed in detail how the burst properties depend on the accretion rate,  composition of the accreted material (the H/He fraction and CNO metallicity), and  internal properties of the  neutron star. For example, for systems accreting mixed H/He, the heat generated from hydrogen burning  is usually a dominant factor for ignition. On the other hand, in evolved  systems in which little or no H is present but the fuel consist mainly of He, the heat required for the burst ignition must come entirely from the electron captures and pycnonuclear reactions in the NS  crust. Comparisons of He-bursts with ignition models thus offer a powerful probe of the physical conditions in the neutron star crust, below the fuel layer, and the cooling processes in the core \\citep{1987ApJ...319..902F,2006ApJ...646..429C}.     The best-known He-accretor is a low-mass binary system 4U~1820$-$30, in which the neutron star orbits its companion once every 11.4 minutes \\citep{1987ApJ...312L..17S}. Such a tiny orbit cannot accommodate a H-rich companion, and the mass donor is likely to be a He-rich white dwarf \\citep{1986Natur.323..105K}. 4U~1820$-$30 is in a bursting mode for around 40 days after switching to the low state \\citep{2001ApJ...563..934C}, while during the rest of its $\\approx$176-day accretion cycle \\citep{1984ApJ...284L..17P} the source does not exhibit bursts \\citep[][and references therein]{1984ApJ...282..713S}. In a 20-hour EXOSAT observation during the low state, \\citet{1987ApJ...314..266H} observed nearly regular bursting from 4U~1820$-$30, detecting  seven bursts with a mean recurrence time of $3.21\\pm0.04$ hours and persistent luminosity  of $L_X=2.8\\times 10^{37}$ ergs s$^{-1}$ between the bursts.      The burst ignition conditions for this source were modelled by \\citet{2003ApJ...595.1077C}, who compared the predicted burst properties with the measurements of \\citet{1987ApJ...314..266H}. \\citet{2003ApJ...595.1077C} presented models for pure He fuel, but also estimated the effect of adding a small amount of hydrogen -- 5$-$35\\% by mass -- as predicted by some stellar evolutionary models \\citep[e.g.,][]{2002ApJ...565.1107P}. The amount of energy released in pycnonuclear and electron capture reactions that escapes from the surface ($Q_{\\rm crust}$) is a free parameter in this model, which also assumes the time-averaged rather than instantaneous accretion rate to set the crust temperature profile, because the thermal time in the crust is much longer than the $\\approx$176-day accretion cycle. \\citet{2003ApJ...595.1077C} found a good agreement between the model, which assumes a mixed fuel (10\\% of hydrogen), and the data, providing that $Q_{\\rm crust}=0.1$ MeV/nucleon  \\citep{2000ApJ...531..988B} and the time-averaged accretion rate is $\\approx$2 times larger than the measured rate. However, for a fuel consisting entirely of He, the required $Q_{\\rm crust}$ was 0.4 MeV/nucleon and the required time-averaged accretion rate was 4--5 times larger.  Self-consistent results were obtained with improved burst ignition models, in which the flux flowing outwards was calculated directly  from the neutron star crust and core neutrino emissivity and the core thermal conductivity \\citep{2006ApJ...646..429C}. \\citet{2006ApJ...646..429C} concluded that, in order to produce the bursts separated by $\\approx$3 hours, the core neutrino emissivity must be very inefficient (e.g., suppresed modified Urca process) and the accretion rate must be $\\approx$2 times larger than that inferred from the X-ray luminosity.          Although all  previous studies of He-bursts have focused on 4U~1820$-$30, the intermittent occurrence of the bursts makes triggering of burst observations  extremely difficult, and only a few recurrence times and corresponding accretion rate estimates are available. A much more suitable candidate for such studies is the source 4U~1728$-$34, which consistently exhibits frequent bursts characteristic of pure He fuel. There is a total of 106 bursts from this source in the {\\sl RXTE} burst catalogue \\citep{2008ApJS..179..360G}. The $\\alpha$ values ($\\approx$200), short rise times and decay time scales suggest a He-rich fuel. The persistent flux during {\\sl RXTE} observations was $1-7 \\times 10^{-9}$ erg cm$^{-2}$ s$^{-1}$ while the burst recurrence times were on average $\\approx$4 hours. Evidence for a short orbital period of 10.77~min has been detected recently in the analysis of {\\sl Chandra} observations (Galloway et al.~2010; in prep.), supporting the long-suspected identification of 4U~1728$-$34  as an ultracompact LMXB. The source is probably accreting pure He from its evolved companion and is practically a twin of 4U~1820$-$30, except for the much more frequent and reliable bursting. 4U~1728$-$34 has also been observed extensively by {\\sl INTEGRAL} \\citep{2006A&A...458...21F,2006AstL...32..456C}. \\citet{2006A&A...458...21F} detected 36 type I bursts during the transition from hard to soft state, where the source luminosity increased from 2$-$12\\% of the Eddington luminosity.       In this paper we present analysis of  new {\\sl Chandra} observations of 4U~1728$-$34. We measured the recurrence times and corresponding accretion rates of the 25 bursts detected  during a 240-ks HETGS exposure. The detailed spectral analysis and the detection of the orbital period and radius expansion bursts are reported in the companion paper (Galloway et al.~2010), while we present the comparison of the observed burst properties  with a  new ignition model. In Section~2 we present the data and describe our analysis of the {\\sl Chandra} observations. The measured burst properties, which include the accretion rates (estimated  from the persistent flux), burst fluences, $\\alpha$ values and burst recurrence times are presented in Section~3. The new burst ignition model is described in Section~4, while Section~5 shows the comparison of the model and data. Possible explanations of the discrepancy between the data and our ignition model are discussed in Section~6. Finally, our conclusions are summarized in Section~7. ",
        "conclusions": "\\label{conclusions}   We compared the observed properties of 38 bursts detected in {\\sl Chandra}, {\\sl RXTE} and  {\\sl INTEGRAL} observations of the helium-rich accretor 4U~1728$-$34 with new ignition models. We find that the observationaly-inferred ignition depths, assuming complete fuel spreading on the neutron star surface, are significantly smaller than the theoretically-derived minimum possible ignition depth of $1-2\\times 10^8$ g cm$^2$.    One way to reconcile the observed and predicted burst recurrence times would be to assume that the observed X-ray luminosities underestimate the accretion rates (for example, due to a non-detection of the  extremely soft or/and hard spectral components). However, for 4U\\,1728$-$34 this scenario is not plausable because these additional spectral components would have to be 2$-$6 times larger than the mesured bolometric luminosity.  In addition, it would imply significantly larger $\\alpha$ values than theoretically predicted.    Alternatively, the ignition column could be increased, without increasing the global accretion rate,  if we assume that the accreted material is confined to a resticted area  on the neutron star surface. To match the observations with our ignition model, we find that the spot radii would have to be in the range 4.5$-$6 km, and that they seem to be correlated with the global accretion rate as predicted by the \\citet{2009ApJ...706..417L}  model. However, detailed ignition models that include all possible effects of the magnetic confinement are needed to confirm this result. An additional confirmation could come from the comparison of the X-ray emission from the accretion disk with the nuclear oscillation  amplitudes detected simultaneously from the source.  To explain the weak, frequent bursts observed from 4U~1728$-$34, we also consider shear-triggered mixing of the accreted helium to larger column depths \\citep{1987ApJ...319..902F,2007ApJ...663.1252P}. However, while this mechanism could push the fuel buffer from the outer envelope closer to the burning zone and explain the observed large $\\alpha$ deviations, it might not be sufficient to explain the observed burst properties alone at the relatively low observed accretion rates.   {\\it We would like to thank Ed Brown, Maurizio Falanga and Stratos Boutloukos for helpful discussions. We would also like to thank the anonymous referee for the useful and constructive comments. }"
    },
    "1005/1005.3938_arXiv.txt": {
        "abstract": "{We show that the extremely low-mass white dwarf NLTT~11748 ($0.17\\,M_\\odot$) is in a close binary with a fainter companion. We obtained a series of radial velocity measurements of the low-mass white dwarf using the H$\\alpha$ core and determined an orbital period of $5.64$ hours. The velocity semi-amplitude ($K=274.8$ km s$^{-1}$) and orbital period imply that it is a degenerate star, and that the minimum mass for the companion is $0.75\\ M_\\odot$ (assuming a mass of $0.167\\,M_\\odot$ for the primary). Our analysis of Balmer line profiles shows that a $0.75\\ M_\\odot$ white dwarf companion does not contribute more than 2\\% or 5\\% of the flux (V-band) for helium- or hydrogen-rich surfaces, respectively. The kinematics of the system suggest that it belongs to the Galactic halo. } ",
        "introduction": "The high proper-motion star NLTT~11748 was identified as an extremely low mass (ELM) white dwarf by \\citet{kaw2009}. The first ELM white dwarfs were discovered to be companions to neutron stars \\citep[e.g.,][]{van1996,bas2003}, but recent follow-up observations of ELM white dwarfs identified in colourimetric or proper motion surveys (e.g., SDSS, NLTT) often find them to be in a close binary with a white dwarf companion \\citep[e.g., SDSS~J125733.63$+$542850.5,][]{bad2009,kul2010,mar2010}. The companions of several other ELM white dwarfs found to be in close binary systems are yet to be formally identified, but they are likely to be more massive white dwarfs \\citep[see,][]{agu2009a,agu2009b}.  The stellar and kinematical properties of NLTT~11748 \\citep{kaw2009} and of two similar objects, LP400-22 \\citep{kil2009,ven2009} and SDSS~J1053+5200 \\citep{kil2010}, also suggest that they are old halo stars. The constraints placed on their total ages help retrace the prior evolution of these systems \\citep[see][]{tau1999,nel2001, nel2004}.  The formation of ELM white dwarfs requires that the systems go through at least two phases of mass transfer, where the second mass transfer phase strips the less massive companion of its outer envelope before the helium ignition. This generic scenario appears valid whether the more massive component is a neutron star or a white dwarf \\citep[see for example][]{nel2001,nel2004}. An example of a possible progenitor to an ELM white dwarf is the bright subluminous B (sdB) star, HD~188112 \\citep{heb2003}. The subdwarf star has a lower mass than average, $0.24$ versus $\\sim0.5\\,M_\\odot$ \\citep{zha2009}, and orbits a degenerate companion with a mass $\\ga 0.73\\,M_\\odot$ every 14.6 hours. The subdwarf is a likely progenitor of a helium white dwarf similar to NLTT~11748. On the other hand, a possible progenitor to an ELM white dwarf with, this time, a neutron star companion is SDSSJ102347.6+003841 \\citep{wan2009}. This is a close binary comprising a G-type star and a 1.69ms pulsar \\citep{arc2009} with an orbital period of 4.75 hours. Although the spectral type of the visible component implies a mass of $\\approx1\\,M_\\odot$, the mass ratio of the system dictates a much lower mass of $\\sim 0.2\\ M_\\odot$ suggesting that the star is significantly evolved with a helium-enriched core. This system probably represents a first link between low-mass X-ray binaries and ELM plus neutron star binaries \\citep{arc2009}.  We report new observations that help clarify the nature of the ELM white dwarf NLTT~11748. We obtained two sets of intermediate resolution spectra that show that NLTT~11748 is in a close binary system (Sect. 2). Our analysis of the spectroscopic data and  a first determination of the binary parameters are presented in Sect. 3. We discuss our results in Sect. 4 and summarise in Sect. 5. ",
        "conclusions": "We find that the ELM white dwarf NLTT~11748 is part of a short-period binary system ($P=5.64$ hrs). Based on our orbital analysis alone, we find that the minimum mass of the companion is $0.75\\,M_\\odot$. However, the companion has independently been identified as a faint white dwarf \\citep{ste2010}, and assuming a mass of $0.75 M_\\odot$, we constrained the companion temperature to $\\le 9600$ K (DA) or $\\le 7200$ K (DC). The kinematics of NLTT~11748 also place the system in the Galactic halo.  \\vspace{-0.2cm}"
    },
    "1005/1005.1536_arXiv.txt": {
        "abstract": "The multifractal nature of solar photospheric magnetic structures are studied using the 2D wavelet transform modulus maxima (\\WTMM{}) method. This relies on computing partition functions from the wavelet transform skeleton defined by the \\WTMM{} method. This skeleton provides an adaptive space-scale partition of the fractal distribution under study, from which one can extract the multifractal singularity spectrum. We describe the implementation of a multiscale image processing segmentation procedure based on the partitioning of the \\WT{} skeleton which allows the disentangling of the information concerning the multifractal properties of active regions from the surrounding quiet-Sun field. The quiet Sun exhibits a average H\\\"older exponent $\\sim -0.75$, with observed multifractal properties due to the supergranular structure. On the other hand, active region multifractal spectra exhibit an average H\\\"older exponent $\\sim 0.38$ similar to those found when studying experimental data from turbulent flows. ",
        "introduction": "\\label{Introduction}  Since the late 70's and the propagation of fractal ideas throughout the scientific community~\\citep{bMan82}, there have been numerous applications of the concepts of scale invariance, self-similarity, long-range dependence in many areas of physics, chemistry, biology, geology, meteorology, economy, social and material sciences~\\citep{b89b,bWes90,b94,bFractal94,b95,b95bis,bFri95,bArn95,bDis02}. Various methods were developed to quantify scale-invariance properties through the computation of the fractal dimension $D_F$ for self-similar objects or the roughness exponent $H$ for self-affine fractals~\\citep{bMan82,b87,bFed88,aArg90,aLea93,b95,aTaq95}. Unfortunately $D_F$ and $H$ are global quantities that do not account for the possibility of point-to-point fluctuations of the scaling properties of a fractal object. The multifractal formalism was introduced in the mid-eighties to provide a statistical description of the fluctuations of regularity of singular measures that are found in chaotic dynamical systems~\\citep{aHas86,aCol87,aRan89} or in modelling of the energy cascading process in turbulent flows~\\citep{aMan74b,aPal87,bMan89,aMen91}. Box-counting and correlation algorithms were successfully adapted to resolve multifractal scaling for isotropic self-similar fractals by computation of the generalized fractal dimensions $D_q$~\\citep{aGras83a,aGras83b,aGra88}. As to self-affine fractals, Parisi and Frisch~\\citep{aPar85} proposed, for the analysis of fully-developed turbulence velocity data, an alternative multifractal description based on the investigation of the scaling behavior of the so-called structure functions~\\citep{bFri95,bMon75}: $S_p(l)=<(\\delta v_l)^p>\\sim l^{\\zeta_p}$ ($p$ integer $> 0$), where $\\delta v_l(x)=v(x+l)-v(x) \\sim l^{h(x)}$ is an increment of the recorded signal over a distance $l$. Then, after reinterpreting the roughness exponent as a local quantity~\\citep{aPar85,aMuz91,aMuz94,aArn95}: $\\delta v_l(x) \\sim l^{h(x)}$ (power-law behavior), the $D(h)$ {\\em singularity spectrum} is defined as the Hausdorff dimension of the set of points $x$ where the local roughness (or H\\\"older) exponent $h(x)$ of $v$ is $h$.  In principle, $D(h)$ can be attained by Legendre transforming the structure function scaling exponents $\\zeta_p$~\\citep{aPar85,aMuz91,aMuz94,aArn95}. Unfortunately, as noticed by \\citep{aMuz91,aMuz93,aMuz94}, both the box-counting and structure function methodology have intrinsic limitations and fail to fully characterize the corresponding singularity spectrum since only the strongest singularities are a priori amenable to these techniques. As such, both methods are limited in their application to real data sets \\citep{aMuz94,Georgoulis:2005,Conlon:2008}.  In previous work, Arneodo and collaborators~\\citep{aMuz91,aMuz93,aMuz94,aArn95} have shown that there exists a natural way of performing a unified multifractal analysis of both singular measures and multi-affine functions, which consists in using the \\textit{continuous wavelet transform}~\\citep{aGou84,aGros84,bMey90,bDau92,bMal98}. By using wavelets instead of boxes, one can take advantage of the freedom of the choice of these ``generalized oscillating boxes'' to get rid of possible smooth behavior that might either mask singularities or perturb the estimation of their strength $h$. The other fundamental advantage of using wavelets is that the skeleton defined by the \\textit{wavelet transform modulus maxima} (\\WTMM{})~\\citep{aMalZho92,aMalHwa92}, provides an adaptative space-scale partitioning from which one can extract the $D(h)$ singularity spectrum \\textit{via} the scaling exponent $\\tau(q)$ of some partition functions defined from the \\WT{} skeleton. The so-called \\WTMM{} method~\\citep{aMuz91,aMuz93,aMuz94,aArn95} therefore gives access to the entire $D(h)$ spectrum \\textit{via} the usual Legendre transform $D(h)=min_q(qh-\\tau(q))$. We refer the reader to \\citet{aBac93} and \\citet{aJaf97b} for rigorous mathematical results and to \\citet{aHen94} for the theoretical treatment of random multifractal functions.  Applications of the \\WTMM{} method to \\UD{} signals have already provided insight into a wide variety of problems~\\citep{refaarn02}, e.g. the validation of the log-normal cascade phenomenology of fully developed turbulence~\\citep{aArn98,aArn99,aDelArn01} and of high-resolution temporal rainfall~\\citep{aVenug06,aVenug06b,aRoux09}, the characterization and the understanding of long-range correlations in DNA sequences~\\citep{aArn95b,aArn96b,aAud01,aAud02}, the demonstration of the existence of a causal cascade of information from large to small scales in financial time series~\\citep{aArn98e,aMuz00}, the use of the multifractal formalism to discriminate between healthy and sick heartbeat dynamics~\\citep{aIva96,aIva99}, the discovery of a Fibonacci structural ordering in \\UD{} cuts of diffusion limited aggregates~\\citep{aArn92,aArn92d,aArn92c,aKuh94} and in hard X-ray emission from solar flares~\\citep{McAteer2007-apj}. The \\WTMM{} method has been generalized from \\UD{} to \\DD{} with the specific goal to achieve multifractal analysis of rough surfaces~\\footnote{The fractal dimension of a rough surface associated to the graph $z=S(x,y)$, where $S(x,y)$ represents the height of the surface at location $(x,y)$, is a quantity $D_F$ between 2 and 3.} with fractal dimension $D_F$ anywhere between 2 and 3~\\citep{aArr97,aArnDec00a,aArnDec00b}. The \\DD{} \\WTMM{} method has been successfully applied to characterize the intermittent nature of cloud structure from satellite images~\\citep{aArn99c,aArnDec00c} and to assist in the diagnosis of breast tissue lesions in digitized mammograms~\\citep{akes01}. In astrophysics, this method was adapted and used to characterize the anisotropic structure of atomic hydrogen gas (HI) in the Galatic disk~\\citep{kha06}. From the analysis of very large mosaics taken from the Canadian Galatic Plane Survey~\\citep{tay03}, directional roughness exponents were introduced to show that the HI in the Galactic spiral arms has a scale-dependent anisotropic signature while the HI in the inter-spiral arm regions exhibits scale-independent anisotropy. Along that line, the \\DD{} \\WTMM{} method was further applied to characterize the space-scale nature of anisotropic structures~\\citep{aSnow08a,aSnow08b,akha09} and to perform objective segmentation of image features of interest from noisy backgrounds~\\citep{kha07,acad07,athi09}. We refer the reader to \\citet{bArn02} for a review of the \\DD{} \\WTMM{} methodology, from the theoretical concepts to experimental applications. Recently, the \\WTMM{} method has been further extended to \\TD{} scalar~\\citep{aKes03_prl} as well as \\TD{} vector~\\citep{akes04prl,aSerra07} fields analysis and applied to \\TD{} (velocity, vorticity, dissipation, enstrophy) numerical data issued from direct numerical simulations (DNS) of incompressible Navier-Stokes equations. Because it combines singular-value decomposition and multifractal description, the so-called tensorial wavelet transform modulus maxima method for vector fields~\\citep{akes04prl,aSerra07} looks very promising for future simultaneous multifractal and structural (vorticity sheets, vorticity filaments) analysis of turbulent flows.  Our aim here is to exploit the ability of the \\WTMM{} method to study compound systems that display some non-analyticity in their multifractal spectra as the signature of some phase transition between two underlying scale invariant components with different multifractal properties~\\citep{bBoh88,aMuz94,aArn95}. These two components can both have some physical significance as previously experienced when using the \\WTMM{} method to detect vorticity filaments in swirling turbulent flows~\\citep{aRouMuz99} or microcalcifications from breast tissue background in digitized mammograms~\\citep{akes01,bArn02}. One of these components can be noise that may cause drastic distortions in the returned multifractal spectra. In this work we will follow a wavelet-based strategy inspired from the one previously used in \\UD{} to detect replication origins and promoters as jumps (discontinuities) in \\UD{} noisy skew profiles in mammalian genomes~\\citep{aBro05,touchon05,nicolay07} and in \\DD{} to perform an objective and automatic segmentation of chromosome territories in fluorescence microscopy imaging of mouse cell nuclei~\\citep{kha07,acad07} and of gold formation on vapodeposited thin gold films~\\citep{athi09}.  The purpose of the manuscript is to demonstrate the suitability and reliability of the \\WTMM{} method to propose a wavelet-based segmentation procedure adapted to solar magnetogram data. In section~\\ref{segmentation}, the basics of the \\DD{} \\WTMM{} method are presented. Its ability to disentangle the underlying scale invariant components of a compound system displaying a phase transition in its singularity spectra is discussed and a strategy of segmentation is implemented. Section~\\ref{section_synthetic_data} is devoted to a test application of the proposed segmentation procedure on a theoretical data set with known multifractal properties. In section~\\ref{observations} we report the results obtained when using this wavelet-based segmentation method to separate active regions from quiet-Sun features in solar line-of-sight magnetogram data. Our conclusion and future directions are then given in Section~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  Many complex physical systems analyzed using fractal and multifractal techniques are surrounded or embedded in a noisy background sometimes originating from instrumental noise. As such these systems are a statistical combination of two distinct self-similar structures. This work addressed the need for an accurate calculation of the multifractal parameters of such complex systems. The presence of compound scale-invariant structures can result in an inaccurate or skewed calculation of the fractal and multifractal parameters when studied as a whole. Using a wavelet-based multi-scale segmentation method, we show that it is possible to disentangle to some extent these two processes and accurately (up to finite-size effects) recover the multifractal characteristics of the system of interest. A theoretical test example for this method was provided in section \\ref{segmentation}. The removal of information relating to the background noise was highlighted in Figure~\\ref{fig_skeleton}. The quantitative results reported in Figure~\\ref{fig_pf} attest of the ability of this segmentation method to recover the multifractal parameters in question.  Let us emphasize here that the application of this method to experimental data for which we do not have a priori knowledge of the possible underlying multifractal processes is a difficult task that requires much attention to perform the most objective segmentation which can not be infered or guided by some physical rule or information. The multifractal analysis is a statistical tool that has direct connections with signal and image processing, but not necessarily with the physics of the system per se. As noticed in section \\ref{sec:sol_mag_active}, the use of a clustering algorithm should greatly help in adjusting the multifractal parameters of the differents components as well as in providing an automated procedure for processing large data sets. This will be reported in a future publication.  The application of this wavelet-based methodology to quiet-Sun data has revealed the multifractal nature of this intermittent noisy component ($<h>=-0.75$) as mainly resulting from the super-granular magnetic structures. The quiet-Sun study was also necessary to get expertise for further analyzing more complex images that involve a segmentation before being able to clearly identify the underlying multifractal properties. We have checked that the partition functions computations for the segmented quiet-Sun phase provide (i) convincing scaling properties and (ii) multifractal spectra $\\tau(q)$ and $D(h)$ estimates in good numerical agreement with the one measured in the previous calibration step.  The assumption of two non-overlapping $D(h)$ is not inconsistent with the data. Let us notice that a phase transition can be observed in the partition function log-log plots when inappropriate segmentation parameters are chosen. In the case of overlapping $D(h)$, there would be a large set of maxima lines in the WT skeleton that could not be genuinely sorted, which would prevent from building accurate multifractal $D(h)$ measurement. From the analyzed data, we were not able to distinguish more than two phases. Finally, this gives us good confidence in the segmentation proposed for solar magnetogram containing an active region. However, further study is needed to precisely quantify scaling properties associated to specific active region features (\\textit{e.g.} emerging magnetic flux along the main polarity inversion line, sunspots build-up, delta-configuration...) and how the \\WTMM{} method can be sensitive to these elements. More precisely, when analyzing higher resolution images than MDI, we expect that this segmentation tool will be all the more necessary as quiet-Sun and active region features are more entangled.   The main outcome of the present study is the demonstration that the proposed multi-scale segmentation procedure provides an objective way of studying the complexity in active regions separately from the surrounding quiet Sun. As such our results are significantly more stable and robust when compared to previous fractal and multifractal analysis~\\citep{Pontieri:2003,McAteer2005-apj}. In a forthcoming paper, we will report on the application of this segmentation method to characterize the evolution of active regions keeping track of the multifractal parameters for possible correlations with extreme solar events~\\citep{Gallagher:2007}. Other applications of this method are in progress as the analysis of the intrinsic multifractal properties of entangled hot and cold interstellar atomic gas from \\TD{} numerical simulations~\\citep{kes_audit:2009}."
    },
    "1005/1005.4914_arXiv.txt": {
        "abstract": "We present simulations of scattering phenomena which are important in pulsar observations, but which are analytically intractable. The simulation code, which has also been used for solar wind and atmospheric scattering problems, is available from the authors. These simulations reveal an unexpectedly important role of dispersion in combination with refraction. We demonstrate the effect of analyzing observations which are shorter than the refractive scale. We examine time-of-arrival fluctuations in detail: showing their correlation with intensity and dispersion measure; providing a heuristic model from which one can estimate their contribution to pulsar timing observations; and showing that much of the effect can be corrected making use of measured intensity and dispersion. Finally, we analyze observations of the millisecond pulsar J0437$-$4715, made with the Parkes radio telescope, that show timing fluctuations which are correlated with intensity. We demonstrate that these timing fluctuations can be corrected, but we find that they are much larger than would be expected from scattering in a homogeneous turbulent plasma with isotropic density fluctuations. We do not have an explanation for these timing fluctuations. ",
        "introduction": "\\label{sec:intro}  Observations of radio pulsars have shown the effects of the interstellar plasma since pulsars were discovered. The first and most obvious effect is dispersion due to the column density of free electrons between the pulsar and the observer. However the effects of scattering due to fluctuations in the electron density were soon recognized as they are very strong in pulsars (Rickett, 1969; Rankin et al., 1970; Cordes, 1986). It is now known that scattering by fluctuations in electron density due to Kolmogorov turbulence are diffractive in nature but the diffractive process is strongly modulated by large scale refraction. As a result there are two spatial scales of the intensity fluctuations, a small scale now called the ``diffractive scale'' $s_\\text{dif}$ and a larger one called the ``refractive scale'' $s_\\text{ref}$ (Prokhorov et al., 1975). The early observations showed only the diffractive scale. The refractive scale was discovered by Sieber (1982) and explained by Rickett et al. (1984). Furthermore the column density changes slowly and this is also observable (Rawley et al., 1988; Ramachandran et al., 2006; You et al., 2007). These fluctuations in ``dispersion measure'' (the pulsar observer's term for column density) have their origin in the same interstellar turbulence and they merge with the diffractive and refractive effects. Observations and theory of interstellar scattering have been reviewed by Rickett (1990) and Narayan (1992).  Scattering is an inherently spatial effect, but it is observed as time variation in, for example, the pulse intensity. The temporal variation is simply caused by spatial variation convected past the observer by the relative motion of the pulsar, the Earth, and the ionized interstellar medium (IISM). Thus we relate an observed time scale $\\tau$ to the corresponding spatial scale S using an effective velocity $S = V_\\text{eff} \\tau$ (Cordes and Rickett, 1998). As there are two spatial scales $s_\\text{dif}$ and $s_\\text{ref}$, there are two temporal scales $\\tau_\\text{dif}$ and $\\tau_\\text{ref}$. In practice $\\tau_\\text{dif}$ is usually minutes to hours so the diffractive component is well sampled in a typical observation. However $\\tau_\\text{ref}$ is usually days to months, so refractive variations are seldom seen in a single observation period. This is why they were not discovered for nearly two decades. Consequently observations are usually analyzed neglecting the refractive variation, but the derived parameters of the diffractive effects are seen to vary with observing epoch (Gupta et al., 1994; Gupta et al., 1999). The question then arises, ``Is this variation part of a continuous variation due to homogeneous turbulence in the IISM, or must we invoke an inhomogeneity in the electron density as was used by Fiedler et al. (1987) to explain extreme scattering events (ESE)?'' ",
        "conclusions": "Simulations have been shown to be a very useful tool in understanding observations of scattering in the interstellar plasma. The simulation engine has been regularized to the point where it can be used by others and it is available from the authors. Here we show a number of examples where simulations have allowed us to understand long standing peculiarities in the observations. We also show that the simulated timing noise due to homogeneous isotropic turbulence is considerably smaller than two cases of observed timing noise which is correlated with intensity fluctuations. This opens a number of interesting questions about the effects of the assumptions of homogeneous isotropic turbulence, and about the observations themselves."
    },
    "1005/1005.3643_arXiv.txt": {
        "abstract": "Solar flares presumably have an impact on the deepest layers of the solar atmosphere and yet the observational evidence for such an impact is scarce. Using ten years of measurements of the Na D$_{1}$ and Na D$_2$ Fraunhofer lines, measured by GOLF onboard SOHO, we show that this photospheric line is indeed affected by flares. The effect of individual flares is hidden by solar oscillations, but a statistical analysis based on conditional averaging reveals a clear signature. Although GOLF can only probe one single wavelength at a time, we show that both wings of the Na line can nevertheless be compared. The varying line asymmetry can be interpreted as an upward plasma motion from the lower solar atmosphere during the peak of the flare, followed by a downward motion. ",
        "introduction": "\\label{sec:intro}    Solar flares are the most powerful events occurring in the solar system. The mechanism invoked for their energy release is the conversion of magnetic energy into radiation, thermal and kinetic energy through the reconnection of magnetic field lines. More precisely, the coronal magnetic field suddenly changes its configuration and free energy when the stress imposed at its footpoints by photospheric plasma motion exceeds a certain threshold \\citep{2005ApJ...635.1299A}. The details of why and how the energy is released are still largely debated (\\textit{e.g} \\citep{2009AdSpR..43..739S}). Once converted, the kinetic energy accelerates particles; part of these are directed downwards and deposit their energy in the solar atmosphere while another part can eventually escape into interplanetary space. The interaction of these accelerated particles with the surrounding plasma ultimately results in electromagnetic radiation at various wavelengths, from X-ray to radio domains \\citep{2008LRSP....5....1B}.  One of the best observed manifestations of flares is the appearance, after the impulsive phase, of loops that emit strongly in the extreme ultraviolet and soft X-rays (SXR). Downward-accelerated particles lose their energy by collisions when they encounter the dense chromospheric plasma. The local material is heated, expands, and then rises to fill the newly configured magnetic-field loop where plasma cools down through radiative and conductive losses. This is the model of chromospheric evaporation \\citep{1984ApJ...287..917A,1987ApJ...317..502F}: the hot plasma located just above the energy-deposition layer evaporates towards the corona. The cool plasma just underneath is pushed down; this is called chromospheric condensation. In this model, the energy deposition layer is located in the middle chromosphere or in the transition region. The rise of chromospheric plasma is manisfested by the observation of Doppler-shifted spectral lines that form at these temperatures \\citep{1987ApJ...317..956S,2005A&A...430..679B} and by the downflow of plasma as observed from the H$\\alpha$ line \\citep{1987SoPh..113..307F}.\\\\   White-light (WL) continuum emission has been also observed in some flares -- that are then called white light flares (WLF) -- but their origin is unclear \\citep{2006SoPh..234...79H}. It is generally agreed that the WL emission takes place in the lower atmosphere, near the minimum temperature region or below \\citep{1986A&A...156...73A,2003A&A...403.1151D}. How these layers are heated, however, remains unknown (see the review by \\citep{2007ASPC..368..417D}). Among the various proposed mechanisms, direct heating by non-thermal electron beams, chromospheric radiative back-warming, or H$^{-}$ continuum emission \\citep{2006ApJ...641.1210X,2000A&A...354..691G} are often cited. When taken individually, none of these mechanisms can provide the amount of energy required by the observations. Direct heating in or near the photosphere has also been suggested \\citep{1994ApJ...429..890D,2009AdSpR..43..739S}. The latter mechanisms suggest that the low solar atmosphere (the chromosphere of course, but also the temperature minimun region and the upper photosphere below) is heated during these flares. As technology evolves, more WL emission is observed during flares \\citep{2008ApJ...688L.119J,Kretzschmar_sub}. We present here observations in visible light made by the \\textit{Global Oscillations at Low Frequency} (GOLF:~Gabriel \\textit{et al}, 1995) instrument onboard SOHO (space mission of international cooperation between ESA and NASA) since 1996 \\citep{1997SoPh..175..207G}. These observations give information on the low solar atmosphere during solar flares.\\\\  One question that we address in this paper is the recurrent problem of retrieving a global picture of the flare since, in addition to the intrinsic variability of the events, the conditions (position, amplitude, ...) vary and the observations are often made by different instruments. We follow here the same approach as in Kretzschmar \\textit{et al}, (2010) and search for a statistical signature of flares by analysing ten years of SOHO/GOLF data consisting in the intensity of the Fraunhofer absorption Na D lines integrated over the full solar disk. These lines are formed in the solar photosphere below the temperature minimum region \\citep{1980A&A....92...63C,1992A&A...265..237B}. We refer here to the photosphere as a few hundred kilometer thick region between the altitude where the optical depth at 500 nm is unity ($ \\tau_{500} = 1 $) and the temperature minimum region.  More precisely, the wings of the Na D line that we study in this paper are mostly generated between 300 and 500 km, just below the transition region and the chromosphere \\citep{2001A&A...371.1128E}. Simulations and observations for a quiet Sun atmosphere have been recently made for the formation of the Na D lines \\citep{2010ApJ...709.1362L}: this study strongly contributes to the idea that these lines are originating from the photosphere. However during a flare, the formation heights of the Na D sodium may differ from quiet conditions. Nevertheless changes in the Na D lines during flares thus indicate that flares affect the lower solar atmosphere including both photosphere and chromosphere. Our analysis therefore provides an opportunity to study the influence of solar flares on the lower solar atmosphere. This is particularly interesting since the best observed flare features usually occur in the upper solar atmosphere, where the contrast between the flare signal and the background emission is largest.\\\\     The intensity of the Sodium line has been measured every five seconds by GOLF. The good cadence, duration and duty cycle of the observations, together with the high quality associated with space-based observations, partly allow us to compensate for the lack of spatial resolution. As we shall see below, large and even medium amplitude flares do have a significant impact on the Na line, which we interpret as the signature of photospheric heating. Furthermore, this work also points to the possible influence of flares on helioseismic studies. \\\\  In Section 2 we introduce the GOLF instrument, the data, and the analysis technique. In Section 3, the results are presented and discussed. Conclusions follow in Section 4. ",
        "conclusions": "The aim of this study was to investigate whether solar flares impact the Na D$_1$ and Na D$_2$ Fraunhofer lines, as measured by GOLF. A statistical analysis based on conditional averaging reveals intensity increases during X and M-class flares, which demonstrates that the Na D lines are affected. Although GOLF observes only one wing at a time, the average evolution of the red and blue wings can be compared in a statistical sense. The observed intensity differences suggest the existence of plasma motions. The blueshift that coincides with the peak of the flare can been interpreted as a upward plasma motion, which is most likely associated with photospheric Na rising towards the chromosphere. The peak of the flare is followed by a redshift, which corresponds to a downward flow. Since the Na D line is mostly generated in the upper photosphere, our study thereby provides strong support for the influence of flares on the lower solar atmosphere.    This interpretation is \\textit{a priori} in contradiction with the classical picture of chromospheric evaporation. It can however be explained if we assume that white-light flares, for which the lower solar atmosphere needs to be heated, have a preponderant effect in our statistics. White-light flares could indeed be more common than expected, and our study supports this idea. However, further quantitative studies are necessary before we can conclude that solar flares do really affect the lower solar atmosphere. Plasma motions could be more accurately determined with more than two measurements on the profile of Na D lines. The GOLF-NG has been proposed for that purpose. Another important conclusion is that future helioseismic instruments should take into account the impact of solar flares in order to improve the measurement of solar oscillations.\\\\     \\textbf"
    },
    "1005/1005.4549_arXiv.txt": {
        "abstract": "{\\small The new release of data from Wilkinson Microwave Anisotropy Probe improves the observational status of relic gravitational waves. The 7-year results enhance the indications of relic gravitational waves in the existing data and change to the better the prospects of confident detection of relic gravitational waves by the currently operating Planck satellite. We apply to WMAP7 data the same methods of analysis that we used earlier [W. Zhao, D. Baskaran, and L.P. Grishchuk, Phys. Rev. D {\\bf 80}, 083005 (2009)] with WMAP5 data. We also revised by the same methods our previous analysis of WMAP3 data. It follows from the examination of consecutive WMAP data releases that the maximum likelihood value of the quadrupole ratio $R$, which characterizes the amount of relic gravitational waves, increases up to $R=0.264$, and the interval separating this value from the point $R=0$ (the hypothesis of no gravitational waves) increases up to a $2\\sigma$ level. The primordial spectra of density perturbations and gravitational waves remain blue in the relevant interval of wavelengths, but the spectral indices increase up to $n_s =1.111$ and $n_t=0.111$. Assuming that the maximum likelihood estimates of the perturbation parameters that we found from WMAP7 data are the true values of the parameters, we find that the signal-to-noise ratio $S/N$ for the detection of relic gravitational waves by the Planck experiment increases up to $S/N=4.04$, even under pessimistic assumptions with regard to residual foreground contamination and instrumental noises. We comment on theoretical frameworks that, in the case of success, will be accepted or decisively rejected by the Planck observations.} ",
        "introduction": "}    The Wilkinson Microwave Anisotropy Probe (WMAP) Collaboration has released the results of 7-year (WMAP7) observations \\cite{wmap7, wmap7-larson}. In this paper, we apply to WMAP7 data the same methods of analysis that we have used before \\cite{stable} in the analysis of WMAP5 data. This is important for updating the present observational status of relic gravitational waves and for making more accurate forecasts for the currently operating Planck mission \\cite{planck}.  In Sec.~\\ref{section2.0} we briefly summarize our basic theoretical foundations and working tools. Part of this material was present in the previous publication \\cite{stable}, but in order to make the paper self-contained we briefly repeat it here. Section \\ref{section2} exposes full details of our maximum likelihood analysis of WMAP7 data. In the focus of attention is the interval of multipoles $2 \\leq \\ell \\leq 100$, where gravitational waves compete with density perturbations. We compare all the results that we derived by exactly the same method from WMAP7, WMAP5 and WMAP3 datasets. This comparison demonstrates the stability of data and data analysis. On the grounds of this comparison, one can say that the perturbation parameters found from the consecutive WMAP data releases have the tendency of saturating near some particular values. The WMAP7 maximum likelihood (ML) value of the quadrupole ratio $R$ is close to previous evaluations of $R$, but increases up to $R=0.264$. The interval separating this ML value from the point $R=0$ (the hypothesis of no gravitational waves) increases up to a $2\\sigma$ level. The primordial spectra remain blue, but the spectral indices in the relevant interval of wavelengths increase up to $n_s = 1.111$ and $n_t=0.111$.    In Sec.~\\ref{section3} we analyze why, to what extent, and in what sense our conclusions with respect to relic gravitational waves differ from those reached by the WMAP Collaboration. The WMAP team has found ``no evidence for tensor modes.\" A particularly important issue, which we discuss in some detail, is the presumed constancy (or simple running) of spectral indices. We derive an exact formula for the spectral index $n_t$ as a function of wavenumbers and discuss in this context the formula for running that was used in WMAP analysis. Another contributing factor to the difference of conclusions is the difference in our treatments of the inflationary ``consistency relations\" based on the inflationary ``classic result.\" We do not use the inflationary theory.    A comprehensive forecast for Planck findings in the area of relic gravitational waves is presented in Sec.~\\ref{section4}. We discuss the efficiency of various information channels, i.e. various correlation functions and their combinations. We perform multipole decomposition of the calculated $S/N$ and discuss physical implications of the detection in various intervals of multipole moments. We stress again that the $B$-mode detection provides the most of $S/N$ only in the conditions of very deep cleaning of foregrounds and relatively small values of $R$. The improvements arising from a 28-month, instead of a nominal 14-month, Planck survey are also discussed. In the center of our attention is the model with the WMAP7 maximum likelihood set of parameters. For this model, the signal-to-noise ratio $S/N$ in the detection of relic gravitational waves by Planck experiment increases up to $S/N=4.04$, even under pessimistic assumptions with regard to residual foreground contamination and instrumental noises. Section \\ref{bayescomp} gives Bayesian comparison of different theoretical frameworks and identifies predictions of $R$ that may be decisively rejected by the Planck observations. ",
        "conclusions": "}      Having analyzed the 7-year data release, the WMAP collaboration concludes that a minimal cosmological model without gravitational waves and with a constant spectral index $n_s$ across the entire interval of considered wavelengths remains a ``remarkably good fit\" to ever improving CMB data and other datasets. The WMAP team emphasizes: ``We do not detect gravitational waves from inflation with 7-year WMAP data, however the upper limits are $16\\%$ lower...\" \\cite{wmap7-larson}, p.11; ``The 7-year WMAP data combined with BAO and ${\\rm H_0}$ excludes the scale-invariant spectrum by more than $3\\sigma$, if we ignore tensor modes (gravitational waves)\" \\cite{wmap7}, p.15; ``We find no evidence for tensor modes...\" \\cite{wmap7}, p.30; ``We find no convincing deviations from the minimal model\" \\cite{wmap7}, p.1, etc.    In contrast, our analysis of WMAP3, WMAP5 and WMAP7 data leads us in the opposite direction: the improving data make the gw indications stronger. The major points of tension between the two approaches seem to be the constancy of spectral indices and the continuing use by the WMAP team of the inflationary theory in data analysis and interpretation. We shall start from the discussion of spectral indices.    The constancy of spectral indices is a reasonable assumption, but not a rule. If the power-law dependence (\\ref{inscf}) is not a good approximation to the gravitational pump field during some interval of time, the constancy of $n_t, ~n_s$ is not a good approximation to the generated primordial spectra (\\ref{primsp}) in the corresponding interval of wavelengths. In fact, the future measurements of frequency dependence of the spectrum of relic gravitational waves will probably be the best way to infer the ``early history of the Hubble parameter\" \\cite{grsol}.    The frequency-dependence of a gw spectrum is fully determined by the time-dependence of the function $\\gamma(t) \\equiv - {\\dot H}/{H^2}$. (In more recent papers of other authors this function is often denoted $\\epsilon(t)$.) The function $\\gamma(t)$ describes the rate of change of the time-dependent Hubble radius $l_{H}(t) \\equiv c/H(t)$: \\[ \\gamma(t) = \\frac{d}{dt}\\left(\\frac{1}{H(t)}\\right) = \\frac{1}{c} \\frac{dl_{H}(t)}{dt}. \\] The function $\\gamma(t)$ is a constant for power-law scale factors (\\ref{inscf}): $\\gamma = (2+\\beta)/(1+\\beta)$, and $\\gamma=0$ for a period of de Sitter expansion. The interval of time $dt$ during the early era when gravitational waves were entering the amplifying regime and their today's frequency spread $d\\nu$ are related by (see Eq.(21) in \\cite{grsol}): \\[ \\frac{d}{dt} =[1-\\gamma(\\nu)]H(\\nu)\\frac{d}{d \\ln \\nu}. \\]  The today's spectral energy density of gravitational waves $\\epsilon(\\nu)$ is related to the early universe parameter $\\gamma(t)$ by (see Eq.(22) in \\cite{grsol}): \\begin{eqnarray} \\label{22} \\gamma(\\nu) = -\\frac{[d~\\ln\\epsilon(\\nu)/d~\\ln\\nu]} {2-[d~\\ln\\epsilon(\\nu)/d~\\ln\\nu]}. \\end{eqnarray}    The spectral index $n_g$ of a pure power-law energy density $\\epsilon(\\nu) \\propto \\nu^{n_g}$ is defined as $n_g = [d~\\ln\\epsilon(\\nu)/d~\\ln\\nu]$. It is reasonable to retain this definition for more complicated spectra. Then, Eq.(\\ref{22}) can be rewritten as \\begin{eqnarray} \\label{ngnu} n_{g}(\\nu) = -\\frac{2 \\gamma(\\nu)}{1- \\gamma(\\nu)}. \\end{eqnarray} Obviously, in the case of pure power-laws (\\ref{inscf}) we return to the constant spectral index $n_g = -2\\gamma/(1-\\gamma) = 2(2+\\beta)$.    Equation (\\ref{ngnu}) was derived for the energy density of relatively high-frequency gravitational waves, $\\nu > 10^{-16} {\\rm Hz}$, which started the adiabatic regime during the radiation-dominated era. In our CMB study we deal with significantly lower frequencies. It is more appropriate to speak about wavenumbers $n$ rather than frequencies $\\nu$, and about power spectra of metric perturbations $h^2(n)$ rather than energy density $\\epsilon(\\nu)$. The $k$-dependent spectral index $n_{t}(k)$ entering primordial spectrum (\\ref{PsPt}) is defined as $n_{t}(k) = [d~\\ln P_{t}(k)/d~\\ln k]$. Then formula for $n_{t}(k)$ retains exactly the same appearance as Eq.(\\ref{ngnu}): \\begin{eqnarray} \\label{ntk} n_{t}(k) = - \\frac{2 \\gamma(k)}{1- \\gamma(k)}. \\end{eqnarray} The spectral index $n_{t}(k)$ reduces to a constant $n_t=2(2+\\beta)$ in the case of power-law functions (\\ref{inscf}).    The spectral index $n_{s}(k) -1$ for density perturbations is defined by $n_{s}(k)-1 = [d~\\ln P_{s}(k)/d~\\ln k]$. The formula for $n_s(k) - 1$ is more complicated than Eq.(\\ref{ntk}) as it contains also $d\\gamma(t)/dt$ as a function of $k$. However, it is important to stress that adjacent intervals of power-law evolution (\\ref{inscf}) with slightly different constants $\\beta$, will result in slightly different pairs of constant indices $n_t,~n_s-1$ in the corresponding adjacent intervals of wavelengths.  Of course, the spectrum itself is continues at the wavelength marking the transition between the two regions.    Extending the minimal model, the WMAP Collaboration works with the power spectrum \\cite{wmap7} \\begin{eqnarray} \\label{runPs} P_{s}(k) = P_{s}(k_0) \\left(\\frac{k}{k_0}\\right)^{n_{s}(k_0) -1 + \\frac{1}{2}\\alpha_s \\ln(k/k_0)}, \\end{eqnarray} which means that the $k$-dependent (running) spectral index $n_{s}(k)$ is assumed to be a constant plus a logarithmic correction: \\begin{eqnarray} \\label{runns} n_{s}(k) = n_{s}(k_0) +\\alpha_s \\ln(k/k_0). \\end{eqnarray} The aim of WMAP data analysis is to find $\\alpha_s$, unless it is postulated from the very beginning, as is done in the central (minimal) model, that $\\alpha_s \\equiv 0$. We note that although logarithmic corrections do arise in simple situations and can even be termed ``natural\" \\cite{grsol}, they are not unique or compulsory, as we illustrated by exact formula (\\ref{ntk}). Nevertheless, we do not debate this point. We accept WMAP's definitions, and we want to illustrate their results graphically, together with our evaluations of $n_s$ in this paper.   The main result of WMAP7 determination is $n_s=0.963 \\pm 0.012$ (68\\% C.L.) derived under the assumption of no gravitational waves and constant $n_s$ throughout all wavelengths included in the considered datasets \\cite{wmap7,wmap7-larson}. When the presence of gravitational waves is allowed, but $n_s$ is still assumed constant, the $n_s$ rises to $n_s= 0.982^{+0.020}_{-0.019}$ from WMAP7 data alone. Finally, from WMAP7 data alone, the WMAP team finds $n_s(k_0)=1.027,~\\alpha_s= -0.034$ in the case of no gw but with running of $n_s$, and $n_s(k_0) =1.076,~\\alpha_s = -0.048$ in the case of running and allowed gravitational waves (we quote only central values without error bars, see Table 7 in \\cite{wmap7}). All the resulting values of $n_s(k)$ derived by WMAP team are shown in Fig.~\\ref{sec2-fig2}. For comparison, we also plot by red lines our evaluations of $n_s$, see Sec.~\\ref{results}.    The lines in Fig.~\\ref{sec2-fig2} show clearly that our finding of a blue shape of the spectrum, i.e. $n_s =1.111$, at longest accessible wavelengths is pretty much in the territory of WMAP findings, if one allows running, even as simple as Eq.(\\ref{runns}), and especially when running is combined with gravitational waves. On the other hand, as was already explained in \\cite{stable}, the attempt of constraining relic gravitational waves by using the data from a huge interval of wavelengths and assuming a constant $n_s$ (or its simple running) across all wavelengths is unwarranted. The high-$\\ell$ CMB data, as well as other datasets at relatively short wavelengths, have nothing to do with relic gravitational waves, and their use is dangerous. As we argued in Sec.~\\ref{results}, the spectral index $n_s$ appears to be sufficiently different even at the span of two adjacent intervals of wavenumbers. The restriction to a relatively small number of multipoles $2 \\leq \\ell \\leq 100$ is accompanied by relatively large uncertainties in $R$, but there is nothing we can do about it to improve the situation, this is in the nature of efforts aimed at measuring $R$. The difference in the treatment of $n_s$ is probably the main reason why we do see indications of relic gravitational waves in the data, whereas the WMAP team does not.    Another contributing factor to the difference of conclusions is the continuing use by WMAP Collaboration of the inflationary theory and its (incorrect) relation $n_t = -r/8$, which automatically sends $r$ to zero when $n_t$ approaches zero. This formula is a part of the inflationary `consistency relations' \\[ r= 16 \\epsilon = - 8 n_t. \\] Only one equality in this formula, $16 \\epsilon = -8n_t$, is correct being an approximate version (for small $\\gamma$, $\\epsilon \\equiv \\gamma$) of our exact formula, Eq. (\\ref{ntk}). The `consistency relation' $r= 16 \\epsilon$ is incorrect. It is an immediate consequence of the ``classic result\" of inflationary theory, namely, the prediction of arbitrarily large amplitudes of density perturbations generated in the limit of de Sitter inflation ($\\epsilon =0,~n_s=1$), regardless of the strength of the generating gravitational field (curvature of space-time) regulated by the Hubble parameter $H$ \\footnote{It is difficult to give adequate references to the origin of the ``classic result\". Judging from publications, conference talks and various interviews, there exists harsh competition among inflationists for its authorship. One popular inflationary activity of present days is calculation of small loop corrections to the theory which is wrong by many orders of magnitude in its lowest tree approximation. For a more detailed criticism of inflationary theory, see \\cite{disca}.}. Certainly, it would be inconsistent, even by the standards of inflationary theory, not to use the relation $r= - 8n_t$ in data analysis, if the inflationary ``classics\" is used in derivation of power spectra and interpretation of results (see, for example, Fig.~19 in \\cite{wmap7}). Obviously, in our analysis we do not use the inflationary theory and its relations.         \\begin{figure} {\\includegraphics[height=10cm]{a8.eps}} \\caption{The spectral index $n_s$ as a function of the wavenumber $k$. The ML results of this work are shown by red lines. Other lines are our plots of WMAP7 findings \\cite{wmap7}, Table 7.} \\label{sec2-fig2} \\end{figure}                 The analysis of the WMAP7 data release amplifies observational indications in favor of relic gravitational waves in the Universe. The WMAP3, WMAP5, and WMAP7 temperature and polarization data in the interval of multipoles $2 \\leq \\ell \\leq 100$ persistently point out to one and the same area in the space of perturbation parameters. It includes a considerable amount of gravitational waves expressed in terms of the parameter $R=0.264$, and somewhat blue primordial spectra with indices $n_s = 1.111$ and $n_t =0.111$. If the maximum likelihood set of parameters that we derived from this analysis is a fair representation of the reality, the relic gravitational waves will be detected more confidently by Planck observations. Even under pessimistic assumption about hindering factors, the expected signal-to-noise ratio should be at the level $S/N = 4.04$.       ~  {\\bf Acknowledgements} We acknowledge the use of the LAMBDA and CAMB Websites. W. Z. is partially supported by Chinese NSF Grants No. 10703005, and No. 10775119 and the Foundation for University Excellent Young Teacher by the Ministry of Zhejiang Education.       \\baselineskip=12truept"
    },
    "1005/1005.4063_arXiv.txt": {
        "abstract": "We study the peculiar velocities of density peaks in the presence of primordial non-Gaussianity.  Rare, high density peaks in the initial density field can be identified with tracers such as galaxies and clusters in the evolved matter distribution. The distribution of relative velocities of peaks is derived in the large-scale limit using two different approaches based on a local biasing scheme.  Both approaches agree, and show that halos still stream with the dark matter locally as well as statistically, i.e. they do not acquire a velocity bias.  Nonetheless, even a moderate degree of (not necessarily local) non-Gaussianity induces a significant skewness ($\\sim 0.1-0.2$) in the relative velocity distribution, making it a potentially interesting probe of non-Gaussianity on intermediate to large scales.  We also study two-point correlations in redshift-space. The well-known Kaiser formula is still a good approximation on large scales, if the Gaussian halo bias is replaced with its (scale-dependent) non-Gaussian generalization.  However, there are additional terms not encompassed by this simple formula which become relevant on smaller scales ($k \\gtrsim 0.01\\iMpch$).  Depending on the allowed level of non-Gaussianity, these could be of relevance for future large spectroscopic surveys. ",
        "introduction": "\\label{sec:intro}  Observations of the large scale structure in the Universe that use galaxies, clusters or other tracers of the density field are done in redshift space: the distance is generally inferred using the redshift $z$, which receives a contribution from the line-of-sight velocity of the object.  These velocities are due to the gravitational field which is correlated with the density field itself.  On large scales where linear perturbation theory in the density field applies, the leading contribution is the squashing (or stretching, in case of an underdensity) of a volume element in redshift space relative to real space.  In this limit, there is a simple relation between the real- and redshift-space power spectra, $P_g$ and $P_{g,s}$, respectively, of a tracer `$g$' \\cite{Kaiser,Fisher95}: \\be P_{g,s}(k, \\mu) = \\left ( 1 + \\frac{f}{b_1} \\mu^2\\right)^2 P_g(k), \\label{eq:Kaiser} \\ee where $\\mu$ is the cosine of the $\\vk$ vector with the line of sight, $f=d\\ln D/d\\ln a$ is the logarithmic derivative of the linear growth factor, and $b_1$ is the linear bias of the tracer population.  Apart from the large-scale, small-correlation limit, the relation \\refeq{Kaiser} makes two assumptions: first, that the tracer population is characterized by a deterministic local bias.  In particular, if the tracer density $\\d_g(\\vx)$ is a local function $F(\\d(\\vx))$ of the matter density perturbation $\\d$, we can expand in $\\d(\\vx)$ to obtain \\cite{FryGaztanaga} \\be \\d_g(\\vx) = b_1\\:\\d(\\vx) + \\frac{b_2}{2} \\d^2(\\vx) +\\dots, \\label{eq:local} \\ee where the bias parameters are either to be seen as free empirical parameters, or can be determined using various theoretical approaches.  Local biasing amounts to the assumption that dark matter halos form in high-density regions (peaks) in the initial density field.  This assumption holds well in the high-peak / massive halo regime, which we assume throughout.  Hence, in the following we will somewhat loosely use ``peaks'' and ``halos'' interchangeably.  The second assumption used for \\refeq{Kaiser} is that the cosmological density field is Gaussian on large scales.  While we will retain the local biasing scheme, we are interested in relaxing the second assumption of Gaussianity. Recently, there has been renewed interest in probing the Gaussianity of the primordial seed perturbations via large scale structure (see \\cite{DesjacquesSeljak10} for a review). The simplest way of obtaining a non-Gaussian field is by adding a \\emph{local} non-linearity: \\be \\Phi(\\vx) = \\Phi_{\\rm G}(\\vx) + \\fNL (\\Phi_{\\rm G}^2(\\vx) - \\<\\Phi_{\\rm G}^2\\>), \\label{eq:philocal} \\ee where $\\Phi_{\\rm G}$ is a Gaussian random field, and $\\Phi$ is the resulting non-Gaussian field. Of course, one can add higher powers to the series \\refeq{philocal}, though the quadratic term usually has the largest impact. Following standard convention, we let $\\Phi$ stand for the primordial potential, related to the density field through the transfer function and Poisson equation (see \\refapp{bispectra}).  As shown by \\cite{DalalEtal08} and confirmed by \\cite{MV08,SlosarEtal},  in the presence of non-Gaussian initial conditions of the local type, halos acquire a scale-dependent correction to their bias which becomes important on large scales: \\be b_1 \\rightarrow b_1 + 2\\fNL (b_1-1)\\d_c \\M(k), \\ee where $\\M(k) \\propto k^{-2}$ is the relation between density and potential in Fourier space (see \\refapp{bispectra}).  Besides the local model, several other possible bispectrum configurations have been proposed in the literature, such as the \\emph{equilateral} (e.g., \\cite{CreminelliEtal06}) and \\emph{folded} types (e.g., \\cite{MeerburgEtal09}).   We summarize these types of non-Gaussianity and the relation between the potential and matter perturbations in \\refapp{bispectra}.  Given the significant impact of (local) non-Gaussianity on the power spectrum of biased tracers, it is then natural to ask what happens to the power spectrum in redshift space, and whether \\refeq{Kaiser} still holds. Furthermore, the distribution of relative velocities between tracers at $\\vx_1$ and $\\vx_2$, \\be \\d \\vu = \\vu(\\vx_2)-\\vu(\\vx_1) \\ee is itself an interesting probe of non-Gaussianity \\cite{CatelanScherrer}.  In this paper, we assume sub-horizon scales throughout, and adopt the Newtonian gauge;  further, we work in Lagrangian coordinates.  At first order, the transition to Eulerian coordinates (which all observations as well as simulation measurements are made in) simply amounts to replacing the Lagrangian bias parameters with their Eulerian counterparts, in case of the linear bias simply $b_{E,1} = b_{L,1}+1$.  While we will focus on the local type of primordial non-Gaussianity for the most part, the expressions obtained can easily be evaluated for any given primordial bispectrum, and we will show selected results for other bispectrum shapes.  The paper is structured as follows: \\refsec{not} introduces density and velocity correlations, and some notation.  \\refsec{deriv} contains two different derivations of the moments of the relative peak velocity distribution. We discuss the distribution of peak velocities in \\refsec{Pdu}. Finally, \\refsec{Pks} presents the power spectrum of peaks (halos) in redshift space.  We conclude in \\refsec{concl}. ",
        "conclusions": "\\label{sec:concl}  In this paper, we have studied the large-scale motions of peaks in the density field in the presence of primordial non-Gaussianity.  In the high-peak regime we are interested in, peaks can be identified with dark matter halos and visible structures such as galaxies and clusters. We have derived the statistics of the relative velocity of halo pairs in two different approaches based on a local biasing scheme.  Note that the bias is local in the physical, \\emph{non-Gaussian} density field. The first, simpler and more accessible approach assumes that on large scales, halo velocities follow those of the total matter.  Given linear and quadratic bias parameters, it is then straightforward to calculate the moments of the relative velocity distribution.  In the second approach (see \\refapp{deriv}), no assumptions are made apart from local biasing and the large-scale, small-correlation limit.  Both approaches agree, showing that the assumption of statistically unbiased halo velocities is consistent on large scales even in the presence of non-Gaussianity.  Interestingly, the presence of non-zero three-point correlations on large scales leads to significant changes in the velocity distribution of halos, and matter itself. For a positive $\\fNL$, the mean streaming velocity is significantly enhanced on large scales, and a non-zero third moment leads to a positive skewness of the velocity difference distribution.  In the local model, the skewness $\\<\\d u^3\\>\\:/\\:\\<\\d u^2\\>^{3/2}$ is of order $(0.1-0.2)\\times(\\fNL/100)$ on a wide range of scales.  Even in the equilateral model, which has little effect on the halo power spectrum, the skewness still reaches 0.06 for the same $\\fNL$. This suggests that the velocity difference distribution can serve as an interesting probe of non-Gaussianity (as already pointed out by \\cite{CatelanScherrer}).  Furthermore, these findings could be of relevance to recent observational reports of significantly higher velocities than expected in the $\\L$CDM (or more generally, General Relativity + Dark Energy) scenario \\cite{FeldmanEtal,KashlinskyEtal,LeeKomatsu}.  In order to evaluate this quantitatively however, one has to take into account non-linear corrections, via perturbation theory and/or simulations.  Our second result is an expression for the redshift-space power spectrum of halos (or matter) in the presence of non-Gaussianity.  The well-known Kaiser formula relating real- and redshift-space power spectra, extended by the scale-dependent halo bias (see e.g. \\cite{DesjacquesSeljak10}), receives non-Gaussian corrections which become relevant on small scales.  Note that through these corrections, the redshift-space power spectrum measured at a given scale as function of line-of-sight angle in principle allows for a direct measurement of non-Gaussianity from a given tracer population, without any reference to the underlying matter power spectrum (in the Kaiser formula at a fixed scale, the non-Gaussian effects are perfectly degenerate with the galaxy bias). Furthermore, the redshift-space power spectrum of matter itself receives corrections of order $\\fNL$ from non-Gaussianity, which are not present in real space.  The non-Gaussian corrections to the Kaiser formula lead to the interesting question of whether there are degeneracies of the effects of non-Gaussianity with other parameters measured from redshift-space distortions, such as dark energy parameters or consistency tests of General Relativity \\cite{ZhangEtal,SongPercival09}.  The severity of the degeneracies will depend on the level of allowed non-Gaussianity.  While the local model is likely to be constrained tightly in the near future, other models such as the equilateral model are much less constrained, but can still be lead to noticeable effects in redshift space (\\reffig{terms-eql}).  Finally, we expect that non-Gaussian effects on halo velocities will have an even larger impact on the \\emph{bispectrum} of halos in redshift space. Again, these questions deserve more study via higher order perturbation theory as well as N-body simulations.   \\vspace*{-0.5cm}"
    },
    "1005/1005.2190_arXiv.txt": {
        "abstract": "We present spectral and kinematic evidence that 2MASS~J06085283$-$2753583 (M8.5$\\gamma$) is a member of the $\\beta$ Pictoris Moving Group (BPMG, age~$\\sim$12~Myr), making it the latest-type known member of this young, nearby association. We confirm low-gravity spectral morphology at both medium and high resolutions in the near-infrared. We present new radial velocity and proper motion measurements and use these to calculate galactic location and space motion consistent with other high-probability members of the BPMG. The predicted mass range consistent with the object's effective temperature, surface gravity, spectral type, and age is 15--35~$M_{Jup}$, placing 2MASS~0608$-$27 well within the brown dwarf mass regime. 2MASS~J06085283$-$2753583 is thus confidently added to the short list of very low mass, intermediate age benchmark objects that inform ongoing searches for the lowest-mass members of nearby young associations. ",
        "introduction": "Low mass members of nearby clusters, moving groups, and loose associations are important objects of study for stellar and planetary astrophysics, formation, and evolution. These groups are older than regions of ongoing star formation (e.g., in Taurus and Orion, ages $\\sim$1--5~Myr) but younger than the disk population (ages $\\sim$1--10~Gyr). Such associations have only recently been discovered because they are low-density groups and their members are spread out on the sky. Currently identified groups include the TW~Hydrae Association (TWA, age $\\sim$10~Myr), the $\\beta$ Pictoris Moving Group (BPMG, $\\sim$12~Myr), and the Tucana/Horologium Association (Tuc-Hor, $\\sim$30~Myr), along with several other clusters/associations with ages younger than $\\sim$500~Myr \\citep[e.g.,][]{Zuckerman04,Torres08}.  Young low mass objects are hotter and have larger radii than older objects of the same mass \\citep[e.g.,][]{Chabrier00}, thus they are brighter and can be studied more efficiently than their older counterparts. Young associations and moving groups are closer to the Sun than the closest star-forming regions ($d$$\\sim$120--140~pc), allowing for more detailed study of their lowest-mass members. Furthermore, low mass members of slightly older associations ($>$20~Myr) have typically ceased accretion and dissipated their primordial disks, making it easier to disentangle atmospheric and circumstellar properties. On the other hand, they still may show signs of youth, such as low gravity spectral features and activity indicators. Eventually a complete census of the low mass members of nearby associations will address outstanding questions in star formation, including the low mass end of the initial mass function and its dependence on environment as well as the mass and age dependence of various processes in star and planet formation.  This Letter establishes 2MASS~J06085283$-$2753583 (hereafter 2MASS~0608$-$27) as the lowest-mass confirmed member of the BPMG. Section~2 describes previous observations of 2MASS~0608$-$27, and Section~3 presents the observations analyzed herein. The kinematic and spatial properties of 2MASS~0608$-$27 are determined in Section~4, and its membership in the BPMG is discussed in Section~5. We present the conclusions of this Letter in Section~6. ",
        "conclusions": "Spectral features, kinematics, and the space location of 2MASS~0608$-$27 clearly confirm it as a member of the BPMG. 2MASS~0608$-$27 is the latest-type non-companion member by a substantial margin. Similarly late-type objects in the BPMG are the M7.5 companion HR~7329B, possibly the M6.5/M9 binary 2MASS~J0041353$-$562112, and the candidate planetary-mass companion to $\\beta$~Pictoris. Ongoing proper motion and RV measurements for apparently young M and L dwarfs in the field will likely confirm larger numbers of very low mass members of the BPMG and other young, nearby loose associations (\\citealt{Schlieder10}, K.~L.~Cruz et al. 2010, in preparation). The identification of the M8.5 object 2MASS~0608$-$27 as a member of the BPMG provides an important late-type benchmark object for further studies."
    },
    "1005/1005.0195_arXiv.txt": {
        "abstract": "A kinetic equation for the collisional evolution of stable, bound, self gravitating and slowly relaxing systems is established, which is valid when the number of constituents is very large. It accounts for the detailed dynamics and self consistent dressing by collective gravitational interaction of the colliding particles, for the system's inhomogeneity and for different constituent's masses. It describes the coupled evolution of collisionally interacting populations, such as stars in a thick disk and the molecular clouds off which they scatter.\\\\ The kinetic equation derives from the BBGKY hierarchy in the limit of weak, but non-vanishing, binary correlations, an approximation which is well justified for large stellar  systems. The evolution of the one-body distribution function is described in action angle space. The collective response is calculated using a biorthogonal basis of pairs of density-potential functions.\\\\ The collision operators are expressed in terms of the collective response function allowed by the existing distribution functions at any given time and involve particles in resonant motion. These equations are shown to satisfy an H-theorem. Because of the inhomogeneous character of the system, the relaxation causes the potential as well as the orbits of the particles to secularly evolve. The changing orbits also cause the angle Fourier coefficients of the basis potentials to change with time. We derive the set of equations which describes this coupled evolution of distribution functions, potential and basis Fourier coefficients for spherically symmetric systems. In the homogeneous limit, which sacrifices  the description of the evolution of the spatial structure of the system but retains the effect of collective gravitational dressing, the kinetic equation reduces to a form similar to the Balescu-Lenard equation of plasma physics. ",
        "introduction": "\\label{introduction}  The description of collisional relaxation in a self-gravitating system usually rests on a Fokker-Planck equation in which the diffusion and braking coefficients are calculated in the local approximation, taking the finite dimension of the system into account by limiting the impact parameter of the collisions to a length of order of the system's size \\citep{Chandra, Chandra43, BinneyTremaine, SpitzerAmas}. Although characteristic relaxation times may be somewhat overestimated by this approximation due to the neglect of collective self-gravitational effects \\citep{Weinberg93}, such a kinetic equation may provide in practice a reasonable description of the collisional relaxation of gravitationally bound systems. It nevertheless rests on assumptions which, from a principle point of view,  are unsatisfactory  because the motion of particles during the collision is regarded as rectilinear and uniform and the system's inhomogeneity, which is basically the reason why collisions with an infinite impact parameter do not occur, is treated by way of an ill-defined cutoff. Moreover, the collective response of the system is not taken into account, since the Fokker-Planck collision term only considers binary collisions between naked particles. A self-gravitating medium, unlike an electrical plasma, does not respond to the presence in it of a particle by screening its interaction potential with other particles. As a result, even distant particles effectively interact, while in electrical, globally neutral, plasmas, the effective interaction distance is limited to the Debye length. In a self-gravitating system, the distance between interacting particles is only limited by the system's inhomogeneity. The spatial structure of the system matters as well as the details of the particle orbits.  \\bigskip  The consistent inclusion of collective screening effects in a kinetic equation for electrically interacting weakly coupled particles has been one of the major theoretical achievements in plasma physics when \\citet{Balescu} and \\citet{Lenard} could derive an equation surpassing in consistency the simple Fokker-Planck equation \\citep{petitSpitzer}. It is the aim of this paper to derive  a similar equation for self-gravitating systems. The task is slightly more difficult because the screening of the electrical interaction at the, usually small, Debye length allows,  in electrically interacting systems, to take the homogeneous and uniform motion limits. These limits cannot be taken in a self-gravitating system. We overcome this difficulty by expressing the kinetic equation in action angle space rather than in position momentum space. This is possible when the Hamiltonian corresponding to the average potential $U({\\mathbf{r}})$ of the system is integrable. It is nevertheless uneasy in general to toggle from one to the other space, although this is certainly possible for spherically symmetric potentials, for flat systems (which may however be unstable) and for special thick disk potentials. Numerical methods could be used to achieve the necessary transformation \\citep{PichonCannon, McMillanBinney}. As an illustrative example, we shall give special attention to spherically symmetric potentials, expanding their kinetic equation into a system which almost entirely avoids any calculation in the position-momentum space. The system's inhomogeneity requires that solutions to the Poisson equation are easily found for any inhomogeneous mass distributions. This is achieved by projecting on a biorthogonal basis of pairs of density-potential functions.  \\bigskip  Many astrophysical systems which have evolved to a quasi-stationary collisionless equilibrium still keep evolving on time scales longer than the dynamical time as a result of gravitational noise induced by their own constituents or by external ones. We disregard external perturbators, which we define as unbound to the system, although, as did \\citet{Weinberg2001II}, these could be treated, if numerous and frequent enough, as a given, non-evolving, population providing a source of gravitational noise for other populations. Loosely bound satellites or remote star populations are regarded as internal to the system. This is possible because our set of kinetic equations allows to simultaneously follow different mass populations. Dwarf satellite galaxies could be regarded for example as one such mass population. Globular clusters, dwarf galaxies, disk galaxies and their haloes are examples of bound systems still evolving as a result of internal noise caused by particle discreteness. Such systems are the object of our study. As in any weakly coupled system, the particles  suffering collisions are dressed by the polarization clouds caused by their own influence on other particles. Collisions between dressed particles have quantitatively different outcomes than collisions between naked ones \\citep{Weinberg98}. This may reflect in significant differences in calculated effective relaxation times and braking or diffusion coefficients, especially when the system, though stable, is not too far from instability \\citep{Weinberg93}. It is therefore useful to account for collective dressing when calculating such processes as secular thick disk evolution, mass segregation in galaxies or in star clusters, or the damping by dynamical friction of galactic populations on high energy orbits. For simplicity, the kinetic equations to be derived below assume that the system is stationary on a dynamical time scale. They thus cannot address questions in which the distribution in angle variable matters, such as the dissolution of freshly accreted satellites, although a simple extension of the theory could. Since however our equations describe the coupled evolution of all populations present in the system, they are well suited to study, for example, the simultaneous evolution by dynamical friction and diffusion of a stellar population and the population of molecular clouds off which these stars scatter.  \\bigskip  The collective response of a self gravitating system to the presence of a perturbing body has been considered by a number of authors, analytically \\citep{Weinberg89, Weinberg95, MuraliTremaine, SahaJog} or numerically \\citep{ThielheimWolff, GnedinOstriker}. Sometimes, the reaction of this perturbation on the perturbing body itself is calculated, as did \\citet{Kalnajs72}, who computed the drag on a large body moving in an homogeneous medium, taking the collective response of this medium into account, and \\citet{WeinbergTremaine}, who considered the global, self-consistent, perturbation caused by a satellite or a barred structure in a spherically symmetric system and its reaction on the perturbator object by the effect of dynamical friction. The secular evolution of the system in response to such perturbations has been considered by \\citet{Weinberg2001}, who considered general types of perturbations on a galaxy, and by \\citet{PichonAubert} who considered perturbations caused by the cosmological environment on dark matter haloes. This evolution is of course in principle observable in N-body simulations, which however have their own difficulties in calculating the long term evolution of such systems \\citep{Binney2004}. A number of authors \\citep{Murali, Weinberg2001, PichonAubert} have studied the collective perturbations caused in a massive spherical galactic halo by its environment. They could calculate the response of this system by resorting to a representation of the particle's motion in action and angle variables, a method first used by \\citet{Kalnajs77}. We follow them on this road. They also made good use of a basis of biorthogonal pairs of density-potential functions. \\citet{Weinberg93} first derived a kinetic equation for the collisional relaxation of a self gravitating system along these lines. His equation accounts for the self-consistent gravitational dressing of the particles, but is otherwise simplified, the geometry supposedly being that of an homogeneously filled periodic cube. The inhomogeneous nature of the system should be described more accurately, still accounting for collective gravitational dressing effects. This is specifically the aim of this paper. \\citet{Chavanis} presented a similar approach to ours for one-dimensional systems, the constituents of which interact by a general long range force. In this paper we further elaborate in section \\ref{secsystemcouple} on the structural evolution of the inhomogeneous system and on the secular evolution of the orbits. ",
        "conclusions": "\\label{secconclu}  Kinetic equations for the collisional evolution of the constituents of self-gravitating inhomogeneous systems have been derived. These equations (\\ref{LAequation}) surpass in consistency the usual Fokker-Planck equations (\\ref{formeFP})~--~(\\ref{Asanseffetscoll})~--~(\\ref{Bsanseffetscoll}). The latter are unsatisfactory from a principle point of view, being local and non-collective.  \\bigskip  By contrast, the proposed equations fully account for the system's inhomogeneity and for the collective gravitational dressing of the colliding particles.  \\bigskip  Equations (\\ref{LAequation}) describe the evolution of distribution functions in action and angle space, which is possible when the hamiltonian associated with the average potential is integrable.  \\bigskip  Physically, these equations describe the evolution of the distribution functions in action space as a result of the weak gravitational noise caused by the discreteness of the particles, dressed with the polarization clouds that their own gravity induces in the system. This gravitational polarization is accounted for in equation (\\ref{LAequation}) in a manner that is fully consistent with the distribution functions, as they are at the moment.  \\subsection{Properties of the kinetic equation} \\label{subsecproperties}  Equation (\\ref{LAequation}) is the sum of a second order derivative term with respect to actions and of a first order one. It therefore basically is of the Fokker-Planck type, although it is definitely simpler in the form of expression (\\ref{LAequation}). The diffusion coefficient involved depends on the 1-body distributions themselves, in particular through the factor $\\mid \\!{\\cal{D}}\\!\\mid^{-2}$ which represents the effect of the dressing of the colliding particles by the gravitational polarization induced around them by their own influence.  \\bigskip  Unlike in electrical plasmas, the polarization dressing in self-gravitational systems does not cause any screening of the interaction, which remains  effective even between distant particles. The mutual distance of such particles is limited only by the finite size of the system. Were the gravitational influence of particles on their surrounding to be neglected, the response matrix $\\varepsilon$ (equation (\\ref{epsilonalphabeta})) would reduce to unity and the coefficients of the corresponding Fokker-Planck kinetic equation would simply be averages by the distribution functions of functions of velocity, as in equations (\\ref{Asanseffetscoll})~--~(\\ref{Bsanseffetscoll}).  \\bigskip  It is apparent from the developments of appendix \\ref{grossesmagouilles}, which lead to equation (\\ref{LAequation}), that the ${\\mathbf{k}}$ component in angle Fourier space of the gravitational polarization response given to a particle has frequency  $\\omega = {\\mathbf{k}}\\! \\cdot \\! {\\mathbf{\\Omega}}$. This means that the polarization cloud which accompanies a particle forms a structure in angle space which vary as ${\\mathbf{w}} - {\\mathbf{\\Omega}}t$: it corotates in angle with that particle.  \\bigskip  The presence of the Dirac function $\\delta({\\mathbf{k}}_1\\!\\cdot  {\\mathbf{\\Omega}}_1 - {\\mathbf{k}}_2\\!\\cdot  {\\mathbf{\\Omega}}_2)$ in equation (\\ref{LAequation}) indicates that particles interact resonantly. This certainly is an important physical property of remote interactions, for which the components of the angle wave vectors ${\\mathbf{k}}_1$ and ${\\mathbf{k}}_2$ must be small. For closer encounters, the modulus of these wave vectors is larger and the resonance condition ${\\mathbf{k}}_1\\cdot  {\\mathbf{\\Omega}}_1 = {\\mathbf{k}}_2\\cdot  {\\mathbf{\\Omega}}_2$ becomes less selective, being more easily satisfied.  \\bigskip  The correlation function has been  calculated on the basis of a linearized theory, which is justified by the weakness of the average interactions in this many-body system. This means that the trajectories of the particles during the collision are regarded as being the unperturbed trajectories. Similarly, the gravitational polarization cloud around any one of the colliding particles is calculated as if the partner in the collision were not present: equation (\\ref{LAequation}) is still a weak collision approximation. A cutoff at small impact parameters is therefore needed to account for the rare strong collisions.  \\bigskip  Equation (\\ref{LAequation}) takes full account of the inhomogeneity  of the system, which is embodied in the dependence of the distribution functions on the actions $\\mathbf{J}$'s. It requires no artificial cutoff at large impact parameters. The details of the trajectories followed by the particles in the present gravitational potential are also fully accounted for, being implicit in  the relations which link the angle and action variables to the position and momentum ones. These relations depend on the actual global gravitational potential of the system, which slowly evolves in time together with the distribution functions.  \\bigskip  The density-potential basis functions $\\psi^{\\alpha}({\\mathbf{r}})$ are choosen at the beginning of the calculation once and for all, but their angle Fourier transforms $\\psi^{\\alpha}_{\\mathbf{k}}({\\mathbf{J}})$, which depend on the actual trajectories of the particles, change with time because the trajectory of a particle of  given actions slowly evolves with the general potential of the system as the relaxation proceeds. As long as it suffers no collision, a given particle keeps its vector  ${\\mathbf{J}}$ fixed because the actions are adiabatic invariants. Collisions, however, cause a secular evolution of the functions $f^a({\\mathbf{J}})$, which is exactly what equation (\\ref{LAequation}) describes.  \\bigskip  \\noindent The description of particle motions is made simple by the use of action and angle variables. Their complexity is embodied in the supposedly known relation between position and momentum variables and action and angle variables. The usefulness of equation (\\ref{LAequation}) is therefore limited to systems for which this relation can be established, either analytically or, possibly, numerically \\citep{PichonCannon, McMillanBinney}.  \\bigskip  \\noindent While the relaxation proceeds, the gravitational potential and the orbits of the particles evolve. As a result, the kinetic equation must be completed by evolution equations for the potential and for other relevant quantities. Section \\ref{secsystemcouple} establishes, for spherical systems, this set of coupled equations.  \\subsection*{Acknowledgements} \\label{acknowledgements}  I am indebted to C. Pichon and D. Aubert for awaking my interest for this problem. Their work \\citep{PichonAubert} introduced me to the tools which make it possible to extend the methods of plasma physics to these systems. I thank the referee, C. Pichon, for many suggestions which helped to improve the first version of this paper. I also thank the Observatoire Astronomique and the Universit\\'e de Strasbourg for accepting me as professor emeritus."
    },
    "1005/1005.4705_arXiv.txt": {
        "abstract": "{% Preliminary results on the discovery and follow-up observations of a new $\\delta$ Scuti pulsator in the Cygnus field are presented. The variability of the star HD 207331 was detected while testing a Str\\\"omgren spectrophotometer attached to the H.~L. Johnson 1.5-m telescope at the San Pedro M\\'artir observatory, M\\'exico.  CCD photometric data acquired soon after confirmed its variability. A few hours of $uvby$ differential photoelectric photometry during three nights revealed at least two beating periods. A two-site observational campaign carried out during one week in 2009 confirms the multi-periodic nature of this new $\\delta$ Scuti pulsator.} ",
        "introduction": "The $\\delta$ Scuti-type pulsators  are stars with masses between 1.5 and 2.5 $M_{\\odot}$ located at the intersection of the classical Cepheid instability strip with the main sequence. They have spectral types A and F, a period range between 0.5 h to 6 h, and generally pulsate with a large number of radial and nonradial modes excited by the $\\kappa$ mechanism. This  makes them interesting targets for seismic studies. Therefore, any new detection of a $\\delta$ Scuti star can be a valuable contribution to asteroseismology.  Since most of the $\\delta$ Scuti stars are short period variables with typical photometric amplitude of 20 mmag, their oscillations can be  easily detected from the ground.  In fact, several $\\delta$ Scuti stars have been discovered accidentally when taken as reference stars of observations of well known $\\delta$ Scuti stars (e.g. Fox Machado et al. 2002 and 2007; Li et al. 2002). Others have been catalogued either in surveys devoted to the characterization of new variables or, as in this case, when considered as constant stars while testing observatory equipment.  In particular, this paper presents a summary of the observations which yielded  the discovery and characterization of the new $\\delta$ Scuti star, HD 207331.   \\begin{figure*} \\includegraphics[width=15cm,height=17cm]{fig1.ps} \\caption{CCD differential light curve:  HD 207331 - Comparison. The fit of the four-frequency solution to the data is shown by solid line.} \\label{fig:curves} \\end{figure*} ",
        "conclusions": "The determination of the evolutionary stage of a field star requires precise estimates of its global parameters. In the case of HD 207331, there is little information in the literature about its physical parameters. In particular, the Hipparcos catalogue (Perryman et al. 1997) provides a parallax of 3.31 $\\pm$ 0.88 mas, from which a distance of 302 pc  can be estimated.  The large relative error of the measured distance ($\\sigma (\\pi))/ \\pi \\sim 0.27$)  implies $\\sigma (M_V)\\sim 0.6$ mag, making the Hipparcos absolute magnitude very imprecise for HD 207331.  Therefore, with the information available it is not possible to make a reliable seismic modeling for HD 207331. However, the complicated oscillation spectrum of HD~207\\-331, with two beating modes, points to it being a fast rotating $\\delta$ Scuti star. This would imply that the four frequencies listed in Table~\\ref{tab:frec} are due to nonradial oscillations. In fact, as is well known, most of the $\\delta$ Scuti stars are rapid rotators with low amplitudes of pulsations, and pulsating with nonradial modes (e.g. Fox Machado et al. 2006).  Nonetheless, more precise information about its evolutionary stage is needed for a more conclusive study.  A summary of the observations which led to the discovery and characterization of the new $\\delta$ Scuti star HD 207331 has been presented. The star shows complicated pulsations as do most of the $\\delta$ Scuti stars.  To date our observations represent the most extensive work on HD 207331.  \\subsection{Future Work} For the future, we will make spectroscopic observations of HD 207331 to obtain an accurate MKK spectral classification to fix more exactly the basic physical parameters of this star in order to be able to make a reliable seismic modeling.  Also, spectroscopic observations will be carried out to test whether this $\\delta$ Scuti star rotates as rapidly as we might expect.  In addition, more differential photometric observations, better distributed in time, are needed to better understand this interesting object."
    },
    "1005/1005.1316_arXiv.txt": {
        "abstract": "% We investigate the use of mid-infrared PAH bands, continuum and emission lines as probes of star-formation and AGN activity in a sample of 100 `normal' and local ($z \\sim 0.1$) galaxies. The MIR spectra were obtained with the {\\it Spitzer} IRS as part of the Spitzer-SDSS-GALEX Spectroscopic Survey (SSGSS) which includes multi-wavelength photometry from the UV to the FIR and optical spectroscopy. The spectra were decomposed using PAHFIT (Smith et al. 2007), which we find to yield PAH equivalent widths (EW) up to $\\sim 30$ times larger than the commonly used spline methods. Based on correlations between PAH, continuum and emission line properties and optically derived physical properties (gas phase metallicity, radiation field hardness), we revisit the diagnostic diagram relating PAH EWs and \\neii/\\oiv\\ and find it more efficient as distinguishing weak AGNs from star-forming galaxies than when spline decompositions are used. The luminosity of individual MIR component (PAH, continuum, Ne and H$_2$ lines) are found to be tightly correlated to the total IR luminosity and can be used to estimate dust attenuation in the UV and in \\ha\\ lines based on energy balance arguments. ",
        "introduction": "We aim at determining the main source of ionizing radiation and star-formation rate of normal galaxies from MIR spectroscopy. ",
        "conclusions": "$\\bullet$ We find systematic trends between MIR spectral properties and optically derived physical properties, in particular between short wavelength PAH EWs and  \\n2ha\\ (gas phase metallicity), and between \\neii/\\oiv\\ versus \\o3hb\\ (radiation field hardness) (Fig. 2, left panel); \\\\ $\\bullet$ The Genzel et al. (1998) diagram has better resolution using PAHFIT than spline decompositions. It is very similar to the optical ``BPT'' diagram (Fig. 2, right panel). The mixed SF/composite region may be revealing obscured AGNs in a large fraction of optically defined ``pure'' SF galaxies.\\\\ $\\bullet$ The PAH, continuum, Ne and H$_2$ luminosities are tightly and nearly linearly correlated to the total IR luminosity, less so to the dust corrected \\ha\\ luminosity (SFR) (Fig. 3, left panel);\\\\ $\\bullet$ Following \\cite{Kennicutt_etal2009}, the MIR components can be used to estimate dust attenuation in \\ha\\  and UV based on energy balance arguments (Fig. 3, right panel)."
    },
    "1005/1005.4019_arXiv.txt": {
        "abstract": "s{ Current analyses of the Lyman-alpha forest assume that the primordial power spectrum of density perturbations obeys a simple power law, a strong theoretical assumption which should be tested. Employing a large suite of numerical simulations which drop this assumption, we reconstruct the shape of the primordial power spectrum using Lyman-alpha data from the Sloan Digital Sky Survey (SDSS). Our method combines a minimally parametric framework with cross-validation, a technique used to avoid over-fitting the data. Future work will involve predictions for the upcoming Baryon Oscillation Sky Survey (BOSS), which will provide new Lyman-alpha data with vastly decreased statistical errors. } ",
        "introduction": "The Lyman-$\\alpha$ forest is the name given to a series of absorption lines in quasar spectra, caused by the scattering of photons via interaction with neutral hydrogen at redshifts $2-4$. At these redshifts, a large proportion of the baryon density of the universe is contained within hydrogen clouds. Most of the hydrogen is ionized, but a small fraction remains neutral, and absorbs photons via the Lyman-$\\alpha$ transition. Hence, the Lyman-$\\alpha$ forest is sensitive to the matter power spectrum on scales from a few up to tens of Mpc, making it the only currently available probe of fluctuations at these weakly non-linear scales. A number of authors have examined the constraints obtainable from the Lyman-$\\alpha$ forest in the past, including Croft et al \\cite{croft}, Gnedin \\& Hamilton \\cite{gnedin}, Viel, Haehnelt \\& Springel \\cite{viel} .  Previous analyses of constraints from the Lyman-$\\alpha$ forest have assumed that the primordial power spectrum is described by a nearly scale-invariant power law. This deserves further attention for a number of reasons. First, it is a strong assumption; if the data are inconsistent with it, derived constraints could be biased to some extent. Second, it is a generic prediction of inflationary models; hence, any test of a power law primordial power spectrum which cannot be attributed to data systematics is a test of inflation. Third, of all current datasets, the Lyman-$\\alpha$ constrains the smallest cosmological scales; thus, it provides the best opportunity presently available to understand the overall shape of the power spectrum. To do this, we shall reconstruct the primordial power spectrum in a minimally parametric way, using a technique called cross-validation to robustly recover the signal. If the data are in agreement with theoretical expectations, the recovered power spectrum will be nearly scale-invariant. In these Proceedings, we discuss a minimally parametric framework for constraining the primordial matter power spectrum, the cross-validation technique, and the methodology for obtaining constraints from observations. Finally, some preliminary results are presented. ",
        "conclusions": ""
    },
    "1005/1005.3852_arXiv.txt": {
        "abstract": "We present a study of the average properties of luminous infrared galaxies detected directly at 24~$\\mu$m in the COSMOS field using a median stacking analysis at 70~$\\mu$m and 160~$\\mu$m.  Over 35000 sources spanning 0$\\leq$$z$$\\leq$3 and $0.06$ mJy$\\leq$$S_{24}$$\\leq$3.0 mJy are stacked, divided into bins of both photometric redshift and 24~$\\mu$m flux.  We find no correlation of $S_{70}/S_{24}$ flux density ratio with $S_{24}$, but find that galaxies with higher $S_{24}$ have a lower $S_{160}/S_{24}$ flux density ratio.  These observed ratios suggest that 24~$\\mu$m selected galaxies have warmer spectral energy distributions (SEDs) at higher mid-IR fluxes, and therefore have a possible higher fraction of active galactic nuclei.   Comparisons of the average $S_{70}/S_{24}$ and $S_{160}/S_{24}$ colors with various empirical templates and theoretical models show that the galaxies detected at 24~$\\mu$m are consistent with ``normal'' star-forming galaxies and warm mid-IR galaxies such as Mrk 231, but inconsistent with heavily obscured galaxies such as Arp 220.  We perform a $\\chi^{2}$ analysis to determine best fit galactic model SEDs and total IR luminosities for each of our bins.  We compare our results to previous methods of estimating $L_{\\rm{IR}}$ and find that previous methods show considerable agreement over the full redshift range, except for the brightest $S_{24}$ sources, where previous methods overpredict the bolometric IR luminosity at high redshift, most likely due to their warmer dust SED.  We present a table that can be used as a more accurate and robust method for estimating bolometric infrared luminosity from 24 $\\mu$m flux densities. ",
        "introduction": "Although rare in the present-day universe, luminous infrared galaxies (LIRGs) were much more numerous in the past, and they may have played a significant role in the evolution of a large fraction of $L>L^{*}$ galaxies \\citep{Sanders:1996p245,Blain:2002p2091,Lagache:2004p2113,LeFloch:2005p2157}.  However, their exact contribution is still poorly understood due to two limitations that have plagued deep surveys performed so far: (1)~the difficulty to identify the most obscured and distant of these objects, as well as measure their redshifts \\citep[they are often faint at optical wavelengths because of dust extinction,][]{Houck:2005p2186}, and (2) the difficulty to accurately characterize their nature (bolometric luminosity, mass, physical processes powering their energy output).  Furthermore, because of limited sensitivity of current space- (\\emph{Spitzer} MIPS) and ground-based (SCUBA, BOLOCAM, MAMBO, AzTEC,...) observations in the far-IR/submillimeter, only a small number of the most luminous of these sources has been studied in detail.  In addition, at high redshifts there are significant limitations due to confusion, which results from the very large instrument beam characterizing current far-IR/submillimeter observations.  Many previous studies of LIRGs have been based on data obtained with the \\emph{Spitzer Space Telescope}, in particular with the Multiband Imaging Photometer \\citep[MIPS,][]{Rieke:2004p2277} at 24~$\\mu$m, the detector's most sensitive band.  Using extrapolations based on libraries of galactic infrared (IR) spectral energy distributions (SEDs), the observed 24~$\\mu$m flux is converted to a bolometric IR luminosity, $L_{\\rm{IR}} \\equiv L(8$--$1000 \\mu$m), which is then used to calculate properties such as instantaneous star formation rate (SFR).  However, at higher redshifts the 24~$\\mu$m band probes shorter rest frame wavelengths, probing rest frame 12 $\\mu$m at $z \\sim 1$, and rest frame 8 $\\mu$m at $z \\sim 2$.  The typical peak of the IR SED of star-forming galaxies and galaxies containing active galactic nuclei (AGN) falls around 50--200 $\\mu$m; at higher redshifts, the 24~$\\mu$m band probes wavelengths farther away from the peak of the IR SED and begins to be heavily affected by broad mid-infrared PAH emission and silicate absorption features.  Observations at longer wavelengths, such as in the \\emph{Spitzer} MIPS  70~$\\mu$m and 160~$\\mu$m bands (which probe rest frame 24~$\\mu$m and 54 $\\mu$m at $z \\sim 2$, respectively), are needed to more accurately characterize the bolometric luminosity, especially at higher redshifts.  However, the MIPS 70~$\\mu$m and 160~$\\mu$m bands are significantly less sensitive and have worse angular resolution than the 24~$\\mu$m band.  This leads to a drastic decrease in the number of sources directly detected at 70~$\\mu$m and 160~$\\mu$m, and the galaxies that are detected are biased toward the most luminous sources.  Therefore, we use a stacking analysis \\citep[as in][]{Dole:2006p1785,Papovich:2007p38} to study the average 70~$\\mu$m and 160~$\\mu$m flux densities of galaxies detected at 24~$\\mu$m.  In using a stacking analysis we lose the ability to study individual galaxies, but find average properties of galaxies that would otherwise be undetectable.  In this work we explore the average mid- to far-IR flux densities of galaxies detected at 24~$\\mu$m and derive a more accurate method to estimate bolometric IR luminosity.  To accomplish this, we measure stacked 70~$\\mu$m and 160~$\\mu$m flux densities of galaxies detected at 24~$\\mu$m in the Cosmic Evolution Survey (COSMOS) field, binned in both redshift and 24~$\\mu$m flux.  We use these stacked fluxes to examine the evolution of mid- to far-IR colors of galaxies as a function of luminosity and redshift.  Our stacked fluxes are fit to libraries of galactic IR SED templates, from which we derive an estimate of the average bolometric IR luminosity.  \\citet{Papovich:2007p38} carried out a similar study employing stacking at 70~$\\mu$m and 160~$\\mu$m, but their analysis was limited by area, with a significantly smaller number of sources.  With an area almost 10 times larger, we obtain more reliable statistics and the ability to bin our sources in narrower bins of redshift and flux.  Our stacked fluxes will eventually be merged with Herschel PACS (100 \\& 160~$\\mu$m), Herschel SPIRE (200--500~$\\mu$m), and SCUBA2 data to get the best sampled SEDs of the high-$z$ literature.  Throughout this work we denote flux density, $f_{\\nu}$ in MIPS 24~$\\mu$m, 70~$\\mu$m, and 160~$\\mu$m bands as $S_{24}$, $S_{70}$, and $S_{160}$, respectively.  When calculating rest-frame quantities, we use a cosmology with $\\Omega_{m} = 0.3$, $\\Lambda = 0.7$, and $H_{0} = 70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\subsection{Comparison to Models}\\label{sec:comparez} In this section, we compare the results of our stacking analysis with the expected fluxes and colors from theoretical models and empirical templates.  We first compare our stacked  $S_{70}/S_{24}$ and $S_{160}/S_{24}$ flux ratios with empirical models of ``normal'' star forming galaxies by \\citet[][hereafter DALE]{Dale:2002p1836} and models of Arp 220 and Mrk 231 from the SWIRE template library \\citep{Polletta:2007p1905}.  We then constrain the average IR luminosity for each bin by fitting to many libraries of theoretical models and empirical templates.  Figure \\ref{fig:comparez} plots $S_{70}/S_{24}$ and $S_{160}/S_{24}$ color as a function of redshift, along with the expected values from the DALE models of star forming galaxies and the SWIRE models of Arp 220 and Mrk 231.  The DALE models are a one parameter family of models and we show models that cover a range of $1 < \\alpha < 2.5$, which describe normal star-forming galaxies with $8.3 < log(L_{\\rm{IR}}) < 14.3$.  Arp 220 is a well-studied galaxy representative of heavily obscured ULIRGs, while Mrk 231 is representative of galaxies with warm mid-IR colors which are known to host AGN.  We use the code \\emph{Le Phare}\\footnote{http://www.cfht.hawaii.edu/~arnouts/lephare.html} developed by S. Arnouts and O. Ilbert to determine the flux ratios of the DALE and SWIRE models at varying redshifts.  \\emph{Le Phare} is a data analysis package used primarily to compute photometric redshifts, but a preliminary phase of the code also computes theoretical magnitudes, given SED libraries and filter bands.  We can see immediately that the average colors determined from our stacking analysis fall within the region spanned by normal star-forming galaxies and Mrk 231, but our colors do not match those of Arp 220 at any redshift.  This suggests that our 24 $\\mu$m selection is biased against heavily obscured objects like Arp 220.  \\subsection{Best-fit Model SEDs and Average Total IR Luminosity}\\label{sec:SEDfit} We use \\emph{Le Phare} to perform a $\\chi^{2}$ analysis to find best-fit galactic model SEDs for the stacked fluxes calculated in each bin.  We use the empirical templates of \\citet{Dale:2002p1836}, \\citet{Lagache:2003p1867},  \\citet{Chary:2001p1425}, and the theoretical radiation pressure models of \\cite{Siebenmorgen:2007p1876}, finding the SED from each library that best fit our stacked data at the correct redshift.   The average $S_{24}$, $S_{70}$, and $S_{160}$ in each bin are plotted along with the best fit models in Figures \\ref{fig:matchz0}--5h, arranged by redshift (Figures 5b--5h are only available in the online version of the paper).  Fluxes from the ``non-detcted'' bins are shown as upper limits.  The parameter space spanned by all four of the models is shaded to give an idea of the spread of possible SEDs that fit the data.  For our model fits, we use four different libraries of ``normal'' star-forming galaxies.  We do have a small fraction of \\emph{XMM-Newton}-detected sources that contain AGN, especially in the highest flux bins, but this does not imply that the mid-IR and/or far-IR fluxes of these galaxies is dominated by the AGN itself (many X-ray sources are PAH dominated in the infrared).  Because of the small number of sources, we are unable to perform a separate stacking analysis of only these X-ray sources.  Although we cannot account for the true contribution of AGN contamination, the tight fits we see suggest that most of our sources are indeed star formation dominated.  From the maximum likelihood function of the $\\chi^{2}$ analysis, we estimate a median $L_{\\rm{IR}}$ and 1$\\sigma$ uncertainties in each bin of our stacking analysis.  We repeat this measurement for each library separately, and then take the mean of the four $L_{\\rm{IR}}$ values to estimate the true luminosity.  The dispersion of the luminosities derived from the different libraries is $\\sim6\\% (3\\sigma)$ in all our bins, which suggests that the four libraries are fairly consistent in their estimates of the best-fit $L_{\\rm{IR}}$.  We add the 1$\\sigma$ uncertainties from each library in quadrature to estimate the error in $L_{\\rm{IR}}$, and find typical $3\\sigma$ errors around 3\\%.  The average IR luminosities  and errors measured for each of our bins are listed in Table \\ref{tab:LIR}.  Bins classified as non-detections at both 70~$\\mu$m and 160~$\\mu$m are marked as upper limits.  As expected, we see that total IR luminosity increases with $S_{24}$ and with redshift.  These galaxies span a large range of IR luminosity, covering ``normal'' galaxies ($L_{\\rm{IR}} \\leq 10^{11} L_{\\sun}$), LIRGs ($10^{11} L_{\\sun} \\leq L_{\\rm{IR}} \\leq 10^{12} L_{\\sun}$), and ultra luminous infrared galaxies (ULIRGs; $L_{\\rm{IR}} \\geq 10^{12} L_{\\sun}$).  To test our uncertainties, we also found best fit SEDs using 70~$\\mu$m and 160~$\\mu$m flux densities that were offset by the errors given in Tables \\ref{tab:70} and \\ref{tab:160} and then measured $L_{\\rm{IR}}$ from these fits.  We find discrepancies much less than 6\\% from this trial, suggesting that our derived infrared luminosities are robust within the errors in our stacking analysis.  Table \\ref{tab:LIR} gives a relation between \\emph{observed} $S_{24}$ and total IR luminosity, with no need for a $k$-correction.  This is an extremely valuable tool given the poor sensitivity of longer wavelength instruments, and provides an effective way to estimate $L_{\\rm{IR}}$ when only having 24~$\\mu$m data.  Some caution must be taken when using Table \\ref{tab:LIR} to estimate $L_{\\rm{IR}}$.  Although the agreement between the best-fit models from many libraries is fairly robust, it is not a complete description of the errors.  For most of these bins, we do not have data on the Rayleigh Jean side of our IR SEDs, which means we cannot estimate the cold dust component and its contribution to $L_{\\rm{IR}}$.  This means that Table \\ref{tab:LIR} is valid under assumption that the libraries used are representative of the diversity of the SEDs beyond $160\\mu$m$/(1+z)$.  These results will be tested in the near future with Herschel and SCUBA2.  Figure \\ref{fig:matchLIR} shows a comparison of $L_{\\rm{IR}}$ calculated from our stacking analysis and $L_{\\rm{IR}}$ calculated from extrapolating mid-IR fluxes from only $S_{24}$ based on the Dale SED libraries, as is common practice.  In general, the two methods are in good agreement at low redshifts, but the $L_{\\rm{IR}}$ calculated using only 24~$\\mu$m data is an overestimate of the true $L_{\\rm{IR}}$ at high redshifts, especially for the brighter $S_{24}$ sources.  This is most likely due to the warmer dust SEDs that we find in the bright $S_{24}$ sources (Section \\ref{sec:color}).  Our results are consistent with the findings of \\citet{Calzetti:2007p1434}, who find that rest-frame 24 $\\mu$m flux is a much better indicator of bolometric infrared luminosity than 8 $\\mu$m flux, and with the findings of \\citet{Papovich:2007p38}, who find that the $L_{\\rm{IR}}$ estimated without taking into account stacked 70~$\\mu$m and 160~$\\mu$m fluxes overestimates the true $L_{\\rm{IR}}$.  To summarize, extrapolating a bolometric infrared luminosity from a 24 $\\mu$m flux density without taking into account 70~$\\mu$m and 160~$\\mu$m flux will result in an overestimate of $L_{\\rm{IR}}$ at high redshifts.  Table \\ref{tab:LIR} will give a more accurate and robust estimate of $L_{\\rm{IR}}$ at these redshifts."
    },
    "1005/1005.3549_arXiv.txt": {
        "abstract": "The photospheres of low-mass red giants show  CNO isotopic abundances that are not satisfactorily accounted for by canonical stellar models. The same is true for the measurements of these isotopes and of the $^{26}$Al/$^{27}$Al ratio in presolar grains of circumstellar origin. Non-convective mixing, occurring during both Red Giant Branch (RGB) and Asymptotic Giant Branch (AGB) stages is the explanation commonly invoked to account for the above evidence. Recently, the need for such mixing phenomena on the AGB was questioned, and chemical anomalies usually attributed to them were suggested to be formed in earlier phases. We have therefore re-calculated extra-mixing effects in low mass stars for both the RGB and AGB stages, in order to verify the above claims. Our results contradict them; we actually confirm that slow transport below the convective envelope occurs also on the AGB. This is required primarily by the oxygen isotopic mix and the $^{26}$Al content of presolar oxide grains. Other pieces of evidence exist, in particular from the isotopic ratios of carbon stars of type N, or C(N), in the Galaxy and in the LMC, as well as of SiC grains of AGB origin. We further show that, when extra-mixing occurs in the RGB phases of population I stars above about 1.2 $M_{\\odot}$, this consumes $^3$He in the envelope, probably preventing the occurrence of thermohaline diffusion on the AGB. Therefore, we argue that other extra-mixing mechanisms should be active in those final evolutionary phases. ",
        "introduction": "Evolved low-mass stars show photospheric CNO isotopic ratios that in many cases are not reproduced by stellar evolutionary codes. In the past years it was recognized by many authors \\citep{boot,wbs,char,sb99} that these chemical anomalies derive from transport mechanisms linking the envelope to zones where partial H-burning occurs. We call these phenomena ``extra-mixing'' or ``deep mixing'' throughout this report. \\citet[][hereafter NBW03]{nol} presented a parametric study of such mixing episodes, suitable to account for the CNO abundances measured in presolar grains of AGB origin. The adopted formalism was based on two parameters, namely the rate of mass transport ($\\dot{M}$) and the temperature ($T_P$) of the deepest zones reached by the circulation. It was also demonstrated that important composition changes can occur without introducing feedbacks on the stellar luminosity, provided $T_P$ is kept low enough (typically, $\\Delta \\log~ T = \\log~T_{\\rm H} - \\log~T_P \\gtrsim$ 0.08$-$0.1, where $T_{\\rm H}$ is the temperature at which the maximum energy of the H-burning shell is released).  Subsequently, physical models for extra-mixing have been explored, which avoid the difficulties previously found with rotationally-induced mechanisms \\citep{pal}. In particular, hydrodynamical models of diffusive processes induced by variations of the mean molecular weight $\\mu$ \\citep{ss69}, called {\\it thermohaline diffusion} \\citep{cz07}, were presented by \\citet{egg1,egg2}. They showed that $^3$He burning into $^4$He and two protons successfully induces the required $\\mu$ inversion ($\\Delta \\mu/\\mu \\simeq -$ 10$^{-4}$), thus driving mixing episodes. Complementarily, \\citet{bwnc}, \\citet{nord} and \\citet{den} suggested that extra-mixing might be driven by magnetic buoyancy, in a dynamo process operating below the envelope.  The model by NBW03 referred explicitly to AGB stars, which are known to be the parents of most presolar grains \\citep{zin1}. Extra-mixing in these objects was considered as necessary by various authors \\citep{hoppe,zinner,heck,nit2}. However, a recent paper \\citep[][hereafter KCS10]{kar} sheds doubts on this requirement, suggesting that slow mass circulation on the RGB might be sufficient. In order to solve the above dilemma, we have to look at the observational data, asking whether there is any clear requirement for extra-mixing on the AGB. Our answer will be that this requirement exists and comes primarily from oxide presolar grains \\citep{nit1,choi,clay,nit2}. Then one must look for a similarly clear constraint from stars, and we find it in the carbon isotope ratios of C(N) giants. Other pieces of evidence (from Ba stars, C-enhanced metal poor stars, etc.), although relevant, will not be addressed here for reasons of space. In order to demonstrate all this we generalize the calculations by NBW03, extending them to cover also RGB stages and including two more values of the initial mass. In Section 2 we compare our results for RGB phases with constraints coming from presolar oxide grains of group 2. In Section 3 we integrate this by considering extra-mixing also in AGB stages; critical tests come again from presolar oxide grains (of group 2, but also of group 1). The evidence from O-rich AGB stars is also discussed. Finally, in Section 4 we comment on the information coming from C(N) stars and we derive preliminary conclusions, also addressing the physical mechanisms necessary to drive the transport. ",
        "conclusions": "Our results indicate that extended mixing on the AGB is required, when considering oxygen isotopes and $^{26}$Al/$^{27}$Al ratios in presolar grains. \\cite[In particular, the oxygen isotopes maintain their crucial role in constraining stellar physics, as early noticed by][]{dearb}. Stellar evidence also points to the same conclusion. This is so in particular for carbon isotopic ratios in C(N) stars. Indeed, while extra-mixing on the RGB is sufficient to explain such ratios (averaging around 60) in the C(N) star sample by \\citet{lamb}, this is no longer true for the stars of \\citet{ab02}. These authors agree rather well (on average) with \\citet{lamb} for the sources in common; however, they also identified a number of new C(N) stars (25\\%) with $^{12}$C/$^{13}$C ratios $\\simeq 10 - 40$. Mainstream SiC grains provide the same evidence (for 23\\% of them, in the St Louis database, the carbon isotope ratio is below 40, while the average value is around 60, as for C(N) stars). \\citet{ab02} found that even assuming, at the beginning of the AGB phase, $^{12}$C/$^{13}$C $\\simeq$ 12, as due to extra-mixing on the RGB, final C-isotope ratios in excess of 43 were always obtained at C/O=1, unless new extra-mixing episodes on the AGB were included. Hence, C(N) giants and SiC grains confirm that extra-mixing occurs on the AGB; this is required also by their oxygen isotopes \\citep{har1,kaha}. The requirements from O-rich stars and grains suggest however a more extended processing than for C(N) stars. As these last have, on average, a larger mass than MS-S giants \\citep{gb08}, the extra-mixing efficiency in population I seems to decrease with increasing stellar mass. For lower metallicities, extra-mixing effects are known to be enhanced \\citep{ch98,grat00}.  In the KCS10 approach, C(N) stars and SiC grains with $^{12}$C/$^{13}$C ratios below about 33 (see also their Table 1) cannot be explained, so that the constraints of SiC grains and carbon stars with low carbon isotope ratios are not reproduced. Note that the difference between the values found by \\citet{kar} and by \\citet{ab02} for the  $^{12}$C/$^{13}$C ratios in C-rich stars (without extra-mixing) is due to different choices for this same ratio at the end of the RGB phase, and for C/O. KCS10 also claim that they ``cover most of the range of observational data points''. They refer to C stars in the Galaxy and in the LMC clusters NGC1846 and NGC1978. Although their Figure 2 might give this impression, this is not the case for the clusters. In fact, O-rich and C-rich stars in each cluster represent an evolutionary sequence, whose C/O and $^{12}$C/$^{13}$C ratios must be explained together. In contrast with this requirement, KCS10 need, for explaining C(N) stars, initial C/O and $^{12}$C/$^{13}$C values of 10 and 0.32; for explaining O-rich AGB stars they instead need 22 and 0.23, respectively (see their Figure 2). KCS10 actually admit their inability at explaining the two constraints together; this failure shows that their approach is not correct, at least for NGC 1846. Conversely, a scenario accounting for both O-rich and C-rich stars was presented by \\citet{leb08}, introducing a moderate extra-mixing on the AGB. These authors also suggested an increased efficiency of extra-mixing for increasing values of the envelope C/O ratio, because the bottom of the convective envelope becomes progressively closer to the H-burning shell while the star climbs the AGB. For the same reason, the extra-mixing efficiency is expected to increase for AGB stars of low metallicity, as the convective envelope becomes hotter for them \\citep{cris}.  The situation of NGC1978 is definitely more puzzling \\citep{lederer}: in that case, a fit to C-rich stars requires the concomitant absence of extra-mixing processes during both the RGB and AGB phases. Thus, it is hard to find a theoretical recipe suitable to reproduce the isotopic ratios of AGB stars in NCG1978 (both O-rich and C-rich) without invoking an ad-hoc solution for this peculiar cluster \\citep[see discussion in][]{lederer}.  \\begin{figure}[t!!] \\centering{\\includegraphics[height=6.0cm]{mixf3.eps}} \\caption{Relative variation of the molecular weight ($\\Delta \\mu /\\mu$) for a 1.5 M$_{\\odot}$ star with a metallicity $Z = Z_{\\odot}/2$, in the layers where $^3$He burns. The inversion, present on the RGB, is reduced or erased on the AGB when $^3$He has been previously consumed (see labels for consumption factors).} \\label{three} \\end{figure}  Finally, we would like to comment on the physical origin of extra-mixing on the AGB. A popular mechanism is today thermohaline diffusion \\citep{egg1,egg2}. However, it is unlikely that it occurs on the AGB of population I stars, when any deep mixing has been previously active on the RGB. Indeed, the envelope abundance of $^3$He, whose burning drives the mixing through a $\\mu$ inversion, would be considerably reduced.  This is shown in Figure \\ref{three}, where we plot the mean molecular weight across the radiative region in RGB and AGB phases, for our 1.5 M$_{\\odot}$ star. While a $\\mu$ inversion of $\\Delta \\mu /\\mu \\simeq -3 \\times 10^{-4}$ (continuous line) is in fact driven by $^3$He burning on the RGB, the AGB cases (dashed lines) are more critical. They refer either to no destruction of $^3$He on the RGB ($^3$He$_{\\rm FDU}$), or to a consumption by factors 3 or 10 (this last case corresponds to the findings by KCS10). It is evident that, for the masses and metallicities considered here, the $\\mu$ inversion is either strongly reduced or erased if $^3$He has been previously consumed. As extra-mixing (hence $^3$He consumption) on the RGB of galactic disk stars is required by observations \\citep{cos,shetr}, the conditions for thermohaline diffusion might be suppressed in their AGB stages \\citep[see also][]{can}. Other mechanisms should therefore be looked for, including magnetic buoyancy \\citep{bwnc}. Note that thermohaline mixing was found to occur on both the RGB and the AGB (first interpulses) in a low-metallicity 1-$M_{\\odot}$ star \\citep{stan}. This is due to the known higher inventory of $^{3}$He in very low-mass stars \\citep{dearb1}.  {\\bf Acknowledgements}. We acknowledge useful comments from G.J. Wasserburg and from two very constructive referee reports. We are indebted to E. Zinner and coworkers for maintaining the on-line repository of presolar grain abundances from which we took the measured data (http://presolar.wustl.edu/~pgd/). C. Abia acknowledges partial support by the Spanish grant AYA2008-04211-C02-02"
    },
    "1005/1005.4475_arXiv.txt": {
        "abstract": "We selected a sample of 24 XMM-Newton light curves  (LCs) of four high energy peaked blazars, PKS 0548$-$322, ON 231, 1ES 1426$+$428 and PKS 2155$-$304. These data comprise continuous light curves of 7.67h to 18.97h in length. We searched for possible quasi-periodic oscillations (QPO) and intra-day variability (IDV) timescales in the LCs of these blazars. We found a likely QPO in one LC of PKS 2155$-$304 which was reported elsewhere (Lachowicz et al.\\ 2009). In the remaining 23 LCs we found  hints of possible weak QPOs in one LC of each of ON 231 and PKS 2155$-$304, but neither is statistically significant. We found IDV timescales that ranged from 15.7 ks to 46.8 ks in 8 LCs. In 13 LCs any variability timescales were longer than the length of the data. Assuming the possible weak QPO periods in the blazars PKS 2155$-$304 and ON 231 are real and are associated with the innermost portions of their accretion disk, we can estimate that their central black hole masses exceed 1.2 $\\times$ 10$^{7}$ M$_{\\odot}$. Emission models for radio-loud active galactic nuclei (AGN) that could explain our results are briefly discussed. ",
        "introduction": "Blazars, including BL Lacertae objects (BL Lacs) and flat spectrum radio quasars (FSRQs), are extragalactic radio sources with relativistic jets aligned nearly ($\\lesssim$ 10$^{\\circ}$) with the line of sight (e.g., Urry \\& Padovani 1995). Blazar emission extends across the entire electromagnetic (EM) spectrum, is predominantly nonthermal, and shows significant polarization in the radio and visible bands where it can be measured.  Blazars show detectable flux variations on diverse timescales ranging from a few  minutes through days and months to decades through all EM bands. Blazar variability timescales have often been somewhat arbitrarily divided into three classes: timescales from minutes to less than a day are called intra-day variability (IDV); those from several days to a few months are known as short timescale variability (STV); while long term variability (LTV) covers changes from several months to many years (e.g., Gupta et al.\\ 2004).  Short and intra-day X-ray variability has been investigated in various classes of AGNs (e.g., Edelson \\& Nandra 1999; Uttley et al.\\ 2002; Markowitz et al.\\ 2003; Vaughan et al.\\ 2003b; McHardy et al.\\ 2004, 2005; Espaillat et al.\\ 2008; Lachowicz et al.\\ 2009; and references therein). The great majority of these studies have been done for the brighter sources, which are typically nearby Seyfert galaxies, and were  reviewed recently by Uttley (2007). There were claims of the detection of QPOs in a few AGNs on the short and IDV timescales in early X-ray observations (Fiore et al.\\ 1989; Papadakis \\& Lawrence 1993; Iwasawa et al.\\ 1998). But all of these claimed QPO detections  were later found to be not statistically significant (Tagliaferri et al.\\  1996; Benlloch et al.\\ 2001; Vaughan 2005; Vaughan \\& Uttley 2006). There have, however, been few stronger  claims of QPO detection recently in X-ray data of various classes of AGNs. The first significant detection of a short X-ray QPO of $\\sim$ 1 hour timescale has been reported for RE J1034$+$396, which is a narrow line Seyfert 1  galaxy (Gierli{\\'n}ski et al.\\ 2008). Espaillat et al.\\ (2008) have reported an X-ray QPO on the timescale of 3.3 ks in 3C 273, a FSRQ. Very recently, Lachowicz et al.\\ (2009) have reported a detection of a X-ray QPO in the BL Lac PKS 2155$-$304 on a timescale of $\\sim$ 4.6 hours. All three of these QPO detections on IDV timescales were based on observations made with XMM-Newton. By using All Sky Monitor data from the Rossi X-ray Timing Explorer, Rani et al.\\ (2009) have reported a probable QPO from the BL Lac  AO 0235$+$164 on a STV timescale of $\\sim$ 18 days.  Such QPOs  may shed new light on the physical processes at the source and associated X-ray emission and so the search for their presence in the light curves (LCs) of AGN is important.  The full spectral energy distributions (SEDs) of blazars have double humped structures with each large spectral feature having comparable total powers. In these double humped SEDs, the first (lower energy) component of the SED is quite clearly dominated by synchrotron radiation from the relativistic jet and the second component is probably due to Inverse-Compton (IC) radiation. On the basis of their SEDs, blazars can be classified into two subclasses known as LBLs (Low-Energy Peaked Blazars) and HBLs (High Energy Peaked Blazars), though those with intermediate peaks are certainly found (Nieppola et al.\\ 2006). In LBLs, the synchrotron component peaks at near-IR, optical or near-UV frequencies and the high energy component usually peaks at GeV energies. In HBLs, the synchrotron component peaks somewhere in the X-ray band and the high energy component is seen to peak (or is extrapolated to do so) at TeV energies.  Here, we study four HBLs which are designated as TeV blazars because of their significant detections in that extreme energy band. Just  six years ago only six TeV blazars were known (see Krawczynski et al.\\ 2004 for a summary of their properties).  Thanks to the development of new TeV facilities such as HESS (High Energy Stereoscopic System; Hofmann et al.\\ 2003; Funk et al.\\ 2004. Aharonian et al.\\ 2006), MAGIC (Major Atmospheric Gamma-ray Imaging Cerenkov; (Baixeras et al.\\ 2004; Cortina et al.\\ 2005; Albert et al.\\ 2006a) and VERITAS (Very Energetic Radiation Imaging Telescope Array System; Holder et al.\\ 2006; Maier \\ 2007, Maier et al.\\ 2008) there has been a complete revolution in TeV gamma-ray astronomy. These groups have detected about a dozen new TeV emitting blazars (HBLs) (Albert et al.\\ 2006b, 2007, 2008, 2009; Acciari et al.\\ 2008, 2009a, 2009b;  Aharonian et al.\\ 2007, 2008a, 2008b, 2008c, 2008d, 2009a, 2009b, 2009c,  and references therein). The enlarged sample of TeV blazars will be very useful for our understanding of the emission mechanism of these extreme blazars through the study of their variability properties across the range of EM bands. Here, we report our search for X-ray variability on IDV timescales in four HBLs from the XMM-Newton satellite public archive. We plan to perform extensive variability studies of HBLs in other EM bands using ground based observations.  The motivation of this work is to examine the nature  of IDV in a sample of four high energy peaked blazars, namely, PKS 0548$-$322, ON 231, 1ES 1426$+$428 and PKS 2155$-$304, in which synchrotron emission from the jets dominates up through the X-ray component. We will  use IDV data to search for QPOs (quasi-periodic oscillations) and variability timescales in the light curves (LCs) of these HBLs, recalling that the shortest timescales can give upper limits to the sizes of the emitting regions. When a blazar is  in a low flux state it is possible that an accretion disk related component of the X-ray emission is dominant and then any IDV timescale could give an estimate of the mass of the black hole presumed to reside at the center of the galaxy (e.g., Gupta et al.\\ 2008, 2009). This possibility arises from the fact that some blazars show evidence of the ``big blue bump'' that is almost certainly indicative of accretion disk emission (e.g., Raiteri et al.\\ 2007) while in their low states, when the jets are weak or even absent.  Under these circumstances the X-ray corona expected above such disks and observed in radio-quiet AGN should be visible.  A break in the power spectral density (PSD) plot from intra-day LCs can also yield the black hole mass of the blazar (e.g., Uttley 2007). XMM-Newton IDV LCs also may possibly be used to make independent estimations of the Doppler factors of the blazars (Fan, Xie \\& Bacon 1999).  The paper is structured as follows.  In Section 2, we give brief descriptions of the data selection criterion and the data reduction method. In Section 3, we discuss the techniques we used to search for variability properties and provide the results in Section 4.  Our discussion is in Section 5 and the conclusions  are given in Section 6. ",
        "conclusions": ""
    },
    "1005/1005.0851_arXiv.txt": {
        "abstract": "Like the majority of spiral galaxies, NGC 6155 exhibits an exponential surface brightness profile that steepens significantly at large radii.  Using the VIRUS-P IFU spectrograph, we have gathered spatially resolved spectra of the system.  Modifying the GANDALF spectral fitting routine for use on the complex stellar populations found in spirals, we find that the average stellar ages increase significantly beyond the profile break radius.  This result is in good agreement with recent simulations that predict the outskirts of disk galaxies are populated through stellar migration.  With the ability to bin multiple fibers, we are able to measure stellar population ages down to $\\mu_V\\sim24$ mag/sq arcsec. ",
        "introduction": "The nature of how and where stellar disks end can shed light on the limits of star formation and the importance of secular evolution.  Stellar disks tend to be embedded in much larger HI disks \\citep{Bosma81,Broeils97,Begum05}, showing that while there are baryons in the outskirts of galaxies, they are not being converted into stars efficiently.  It is often assumed that there is a surface density threshold for star formation \\citep{Kennicutt89} which would provide a natural mechanism for truncating the stellar component of galaxies.   The surface brightness profiles of spiral galaxies are commonly fit with exponentials \\citep{Freeman70}.  While the exponential function is a convenient description, few astronomers would suggest that stellar disks extend to infinite radii.  Observations of edge-on systems originally suggested that stellar disks have well defined truncation radii \\citep{vdk81a,Kregel04}.  Using large samples of nearly face-on galaxies, \\citet{Pohlen06} found a variety of behaviors at large galactic radii.  \\citet{Pohlen06} find that only 10\\% of disks are well fit by a single exponential light profile, while in 60\\% of their sample the surface brightness profile is well fit by a downbending broken exponental, with a break radius 1.5-4.5 times the inner scale length.  The final 30\\% of their sample are described with a broken upbending exponential profile.  Unlike the edge-on studies, they do not find evidence for sharp truncations.    \\citet{Roskar08} present a possible formation mechanism for surface brightness profile breaks.  They find that a star formation threshold seeds the onset of a truncation, but radial migration populates the disk beyond the break producing a down-bending broken profile.  Since the stellar migration is a random walk, only ancient stars have had enough time to get to very large radii beyond where star formation shuts off.  These simulations thus predict that the break radius should correspond with an increase in the average stellar ages.  The formation and appearance of outer disks has since been investigated in cosmological simulations by several groups who also find an upturn in age at the break radius \\citep{Sanchez09b,Martinez09}. However, the relative importance of radial migration and in-situ star formation in a fully-cosmological setting remains an open question.    While simulations find that profile breaks can be caused by a combination of a star formation threshold and stellar migration, observations suggest that star formation can proceed even in the outskirts of disks.  Deep \\ha\\ imaging by \\citet{Ferguson98} show signs of extended star formation.  \\citet{Thilker08} also find that $\\sim30$\\% of spiral galaxies have UV emission extending beyond their traditional star formation threshold radius.  Similarly, \\citet{Herbertfort09} find clustering of point sources around NGC 3184, implying star cluster formation in the outer disk.   Recent observations of resolved stars have found a variety of results for stellar truncation regions.  \\citet{Williams09} find that in low-mass M33 the the age gradient changes from negative to positive at the truncation radius.  Meanwhile, \\citet{deJong07} find that the truncation region in NGC 4244 is the same for stars of all ages.  \\citet{deJong07} interpret this as a sign the truncation is the result of a dynamical interaction, but these results were subsequently shown to be consistent with the radial migration model \\citep{Roskar08}.  See \\citet{Vlajic10} for a review of the recent advances in studying the outskirts of spiral disks.  The observational results point to an interesting conundrum that while many studies find active star formation in the outskirts of disk galaxies, the stars in the region are often older than the inner disk.   To date, truncation regions have been studied with broadband colors and resolved stars.  There are a limited number of systems we can resolve, and faint broadband surface photometry is dominated by systematic errors associated with flat-fielding and sky subtraction.  While there have been searches targeting emission lines \\citep{Ferg98,Chirst10}, in this letter we present the first spectroscopic study of stellar continuum spectral features across a surface brightness profile break. ",
        "conclusions": "\\citet{Pohlen06} describe NGC 6155 as lopsided and classify the galaxy as a Type II-AB (a downbending profile in an asymmetric disk).  While our fiber-photometry shows slight lopsidedness, the velocity field is quite regular and there are no nearby neighbors.  All the galaxies within a 1 Mpc projected radius and similar redshift of NGC 6155 are at least 3 magnitudes fainter.  It therefore seems unlikely that the profile break in NGC 6155 is the result of a dynamical interaction with a neighboring galaxy.  Our result that average stellar ages increase beyond the break radius is consistent with the results of \\citet{Bakos08} who stacked multiple SDSS images and found the $g-r$ color became redder beyond the profile break.  This is also similar to the age gradient observed in M33, with the age decreasing out to the profile break where the gradient reverses and the stellar population becomes older beyond the break \\citep{Williams09}.  \\citet{Bakos08} claim that while their stacked surface brightness profiles shows a break, the stellar mass profile is described by a single exponential.  Our model mass-to-light values are not well constrained beyond the profile break, however, the increase in age we detect does suggest the M/L is larger in this region.  Therefore we expect the stellar mass profile to be smoother than the light profile.   The increase in stellar age we measure is very similar to the simulations of \\citet{Roskar08} where the average stellar age jumps from 4 to 6 Gyr.  The best fits beyond the break radius have $\\tau\\sim0-9.5$ Gyr, consistent with significant ongoing star formation in the region.  We therefore take our results as evidence that while stellar migration can populate outer disks, our fits are consistent with some low-level \\emph{in situ} star formation as well.  Along with the change in stellar age, the surface brightness break also corresponds with a lack of line emission (Figure~\\ref{sb}, lower left).  This could be a simple coincidence where the break radius is also where the emission becomes faint enough that we no longer detect it, but it is also consistent with the theory that the break radius represents the location of a star formation threshold.  It should be pointed out that our observations are also consistent with a radial change in the stellar IMF.  If the outer disk only hosts lower-mass molecular clouds, we could expect star formation to be unable to produce the highest mass stars \\citep{Kroupa03,Koeppen06}.  Because young massive stars contribute heavily to the overall luminosity, their absence would result in a young stellar population that spectroscopically appears older.  A change in the IMF can also explain the lack of emission lines beyond the break \\citep{Pflamm09}.   In a forthcoming paper, we will present a larger sample of similar observations.  Our larger sample includes galaxies with no surface brightness profile break as well as up-bending profiles.  We also plan to further refine our spectral synthesis template fitting method to improve the metallicity measurements of the stellar populations."
    },
    "1005/1005.0256_arXiv.txt": {
        "abstract": "We measure the topology of the main galaxy distribution using the Seventh Data Release of the Sloan Digital Sky Survey, examining the dependence of galaxy clustering topology on galaxy properties. The observational results are used to test galaxy formation models.  A volume-limited sample defined by $M_r<-20.19$ enables us to measure the genus curve with amplitude of $G=378$ at $6h^{-1}$Mpc smoothing scale, with 4.8\\% uncertainty including all systematics and cosmic variance. The clustering topology over the smoothing length interval from 6 to $10 h^{-1}$Mpc reveals a mild scale-dependence for the shift ($\\Delta\\nu$) and void abundance ($A_V$) parameters of the genus curve. We find substantial bias in the topology of galaxy clustering with respect to the predicted topology of the matter distribution, which varies with luminosity, morphology, color, and the smoothing scale of the density field. The distribution of relatively brighter galaxies shows a greater prevalence of isolated clusters and more percolated voids. Even though early (late)-type galaxies show topology similar to that of red (blue) galaxies, the morphology dependence of topology is not identical to the color dependence.  In particular, the void abundance parameter $A_V$ depends on morphology more strongly than on color. We test five galaxy assignment schemes applied to cosmological N-body simulations of a $\\Lambda$CDM universe to generate mock galaxies: the Halo-Galaxy one-to-one Correspondence model, the Halo Occupation Distribution model, and three implementations of Semi-Analytic Models (SAMs).  None of the models reproduces all aspects of the observed clustering topology; the deviations vary from one model to another but include statistically significant discrepancies in the abundance of isolated voids or isolated clusters and the amplitude and overall shift of the genus curve.  SAM predictions of the topology color-dependence are usually correct in sign but incorrect in magnitude. Our topology tests indicate that, in these models, voids should be emptier and more connected, and the threshold for galaxy formation should be at lower densities. ",
        "introduction": "Galaxy clustering has long been used to constrain cosmological models and to understand formation of galaxies. The most extensively studied clustering statistics are the autocorrelation function and the power spectrum. They are Fourier transforms of each other, and measure the clustering strength as a function of scale. The ``Standard Cold Dark Matter'' model (with the density parameter $\\Omega_m=1$, scale-invariant primordial fluctuations, and standard relativistic particle background), popular in the 1980s, was ruled out by showing that it was inconsistent with the observed galaxy correlation function (Maddox et al. 1990) and power spectrum (Vogeley et al. 1992; Park et al. 1992, 1994). These statistics are also used to measure the biasing in the galaxy clustering amplitude with respect to matter, and are key constraints on galaxy formation models connecting dark matter halos and luminous galaxies such as semi-analytic models (SAM; e.g., White \\& Frenk 1991; Kauffmann et al.\\ 1993; Cole et al. 1994, 2000; Benson et al. 2002; Bower et al. 2006; Cattaneo et al. 2006; Croton et al. 2006; De Lucia \\& Blaizot 2007; Monaco et al. 2007; Somerville et al. 2008) and halo occupation distribution models (e.g., Seljak 2000; Peacock \\& Smith 2000; Scoccimarro et al. 2001; Berlind \\& Weinberg 2002; Kang et al. 2002; Berlind et al. 2003; Zheng, Coil, \\& Zehavi 2007).  Topology analysis was introduced by Gott et al. (1986, 1987) to test the Gaussianity of the primordial density fluctuations, which is one of the key characteristics of simple inflationary models (Bardeen et al. 1986). At large scales density fluctuations are still in the linear regime and maintain their initial topology, and it is possible to check whether or not the primordial fluctuations were a Gaussian field. At smaller scales, non-linear gravitational evolution and biased galaxy formation make the topology of the observed galaxy distribution deviate from the Gaussian form even if the initial conditions were Gaussian distributed as shown by perturbation theories and large N-body simulations (Park \\& Gott 1991; Weinberg \\& Cole 1992; Matsubara 1994; Matsubara \\& Suto 1996); using fractional volume rather than density threshold as the independent variable in topology analysis mitigates but does not eliminate these non-linear and biasing effects (Weinberg et al.\\ 1987; Melott et al.\\ 1988). Through studies of many observational samples, the topological properties of the large-scale distribution of galaxies have been examined (Gott et al. 1989; Park, Gott \\& da Costa 1992; Moore et al. 1992; Vogeley et al. 1994; Rhoads, Gott, \\& Postman 1994; Protogeros \\& Weinberg 1997; Canavezes et al. 1998; Park, Gott \\& Choi 2001; Hoyle, Vogeley \\& Gott 2002; Hikage et al. 2002, 2003; Park et al. 2005; James, Lewis \\& Colless 2007; Gott, Choi \\& Park 2009; James et al. 2009). On non-linear or quasi-linear scales, topology analysis is useful in constraining both cosmological parameters and galaxy formation mechanisms (Park, Kim \\& Gott 2005; Gott et al. 2008). In particular, differences in clustering topology for different types of galaxies reflect their different history of formation and evolution. Therefore, looking at the topology of large-scale structure traced by different types of galaxies can put strong constraints on galaxy formation mechanisms.  In this paper we use the Seventh Data Release (DR7; Abazajian et al. 2009) of the Sloan Digital Sky Survey (York et al.\\ 2000) to measure the topology of the galaxy distribution and its dependence on galaxy luminosity, color, and morphology.  DR7 constitutes the final release of the SDSS Legacy Survey, and thus of the main SDSS galaxy redshift survey.  We supplement the SDSS data with missing redshifts to increase the completeness of the redshift catalog. We then generate a set of volume-limited samples of the SDSS galaxies divided according to their luminosity, morphology, and color to study the relation between the topology and properties of galaxies tracing the large-scale structure. ",
        "conclusions": "We use the SDSS DR7 main galaxy catalog supplemented with missing redshifts and with increased spectroscopic completeness to measure the galaxy clustering topology over a range of smoothing scales. The distribution of galaxies observed by the SDSS reveals extremely diverse structures.  A volume-limited sample, BEST defined by $M_r<-20.19$, enables us to measure the genus curve with amplitude of $G=378\\pm 18$ at a smoothing scale of $6h^{-1}$Mpc, with the quoted uncertainty including all systematics and cosmic variance. The amplitude is 5.4 times larger than our previous measurement using the SDSS DR3 sample (Park et al. 2005) and the uncertainty decreases from 10.5\\% to 4.8\\% at the same smoothing length of $R_G=6h^{-1}$Mpc. We calculate the galaxy clustering topology over the interval from $R_G=6h^{-1}$Mpc to $10h^{-1}$Mpc, and find mild scale-dependence for the shift ($\\Delta\\nu$) and void abundance ($A_V$) parameters. The measured genus curve is qualitatively similar to the form predicted for Gaussian primordial fluctuations (Hamilton et al. 1986), but the differences are statistically significant at these scales: a shift of the peak towards negative $\\nu$, and fewer isolated voids and isolated clusters than the Gaussian prediction ($A_C$ and $A_V < 1$). The bias in topology of galaxy clustering with respect to that of matter is measured by assuming that the matter density field is given by our $\\Lambda$CDM N-body simulation. We detect strong topology bias in galaxy clustering, which is also scale-dependent.  We confirm the luminosity dependence of galaxy clustering topology discovered by Park et al. (2005). The distribution of brighter galaxies is more shifted towards ``meatball' topology (lower $\\Delta\\nu$) and shows greater percolation of voids (lower $A_V$). We find galaxy clustering topology depends also on morphology and color. Even though early (late)-type galaxies show topology similar to that of red (blue) galaxies, morphology-dependence of topology is not identical with color-dependence. In particular, the void abundance parameter $A_V$ depends on morphology more strongly than color.  We tested five galaxy formation models, which are used to assign galaxies to the outputs of N-body simulations.  Three of these are semi-analytic models, one is an HOD model that assigns galaxies to halos with parameters tuned to match other clustering statistics, and one is a scheme that assigns galaxies to halos and subhalos. None of them reproduces all aspects of the observed topology, though the differences from one model to another are comparable to the discrepancies with the observations. For this reason, and because the initially Gaussian $\\Lambda$CDM model successfully reproduce the observed topology of LRGs at large scales, we attribute the discrepancies to failures of the galaxy formation model rather than non-Gaussian initial conditions. The semi-analytic models can also predict the topology of color subsets, but none of them fully captures the observed topology differences between red and blue galaxies.  In future work, we will investigate models with non-Gaussian initial conditions to see what levels of primordial non-Gaussianity can be ruled out by our measurements.  In principle, the high-precision topology measurements presented here and by Gott et al. (2009) can provide valuable constraints on non-standard inflationary models or alternative hypotheses for the origin of primordial fluctuations.  Appendix A details our estimates of systematic biases in the genus curve measurements, demonstrating that the dominant effect is peculiar velocity distortions in redshift space rather than sample geometry or boundary effects.  Except where noted otherwise, observational measurements in this paper are corrected for these biases, so they can be compared to theoretical predictions in real space with periodic boundaries.  Appendix B investigates error covariances, showing that while the individual points on the genus curve have strongly covariant errors, the statistics $G$, $\\Delta\\nu$, $A_V$, and $A_C$ are approximately independent.  Appendix C tabulates the full genus curves for our best samples, complementing the statistics recorded in earlier tables.  While future surveys will use luminous galaxies and emission-line galaxies to probe structure in the distant universe, the SDSS DR7 sample is likely to remain the definitive map of large scale structure at low redshift traced by a broad spectrum of galaxy types, for the foreseeable future.  The measurements in this paper characterize the topology of this definitive sample, attaining unprecedented statistical precision and providing a valuable test for future models of primordial fluctuations and galaxy formation physics."
    },
    "1005/1005.5645_arXiv.txt": {
        "abstract": "The fundamental plane of early-type galaxies is a rather tight three-parameter correlation discovered more than twenty years ago. It has resisted a both global and precise physical interpretation despite a consequent number of works, observational, theoretical or using numerical simulations. It appears that its precise properties depend on the population of galaxies in study. Instead of selecting a priori these populations, we propose to objectively construct homologous populations from multivariate analyses. We have undertaken multivariate cluster and cladistic analyses of a sample of 56 low-redshift galaxy clusters containing 699 early-type galaxies, using four parameters: effective radius, velocity dispersion, surface brightness averaged over effective radius, and Mg2 index. All our analyses are consistent with seven groups that define separate regions on the global fundamental plane, not across its thickness. In fact, each group shows its own fundamental plane, which is more loosely defined for less diversified groups. We conclude that the global fundamental plane is not a bent surface, but made of a collection of several groups characterizing several fundamental planes with different thicknesses and orientations in the parameter space. Our diversification scenario probably indicates that the level of diversity is linked to the number and the nature of  transforming events and that the fundamental plane is the result of several transforming events. We also show that our classification, not the fundamental planes, is universal within our redshift range (0.007 -- 0.053). We find that the three groups with the thinnest fundamental planes presumably formed through dissipative (wet) mergers. In one of them, this(ese) merger(s) must have been quite ancient because of the relatively low metallicity of its galaxies, Two of these groups have subsequently undergone dry mergers to increase their masses. In the k-space, the third one clearly occupies the region where bulges (of lenticular or spiral galaxies) lie and might also have formed through minor mergers and accretions. The two least diversified groups probably did not form by major mergers and must have been strongly affected by interactions, some of the gas in the objects of one of these groups having possibly been swept out. The interpretation, based on specific assembly histories of galaxies of our seven groups, shows that they are truly homologous. They were obtained directly from several observables, thus independently of any a priori classification. The diversification scenario relating these groups does not depend on models or numerical simulations, but is objectively provided by the cladistic analysis. Consequently, our classification is more easily compared to models and numerical simulations, and our work can be readily repeated with additional observables. ",
        "introduction": "Physical understanding of astrophysical objects most often uses correlation diagrams. For early-type galaxies, such scaling laws have been for instance established on one hand between optical luminosity and central velocity dispersion $\\sigma$ \\citep{FaberJackson1976}, and on the other hand between their effective radius $R_e$ and surface brightness averaged over effective radius $<\\mu_e>$ \\citep{Kormendy1977}. These correlations are rather tight, but the scatter is still reduced using a three-parameter relation of the form : $\\log{R_e} = a \\log\\sigma + b <\\mu_e> +\\ c$ \\citep{Dressler1987,Djorgovski1987}. This relation extends to faint and low-mass galaxies \\citep[e.g.][]{Nieto1990}. This is the fundamental plane (hereafter FP) of early-type galaxies.  A long-running difficulty is the so-called ``tilt'' of the FP with respect to the ``virial plane'' obtained with the virial theorem and some simple assumptions about the population of early-type galaxies. Indeed, this tilt is different for different types of galaxies, like disk ones \\citep[e.g.][]{Robertson2006,Hopkins2008}. Many studies have been devoted to this problem without a clear solution. The motivation is to obtain pure correlations both to constrain the models better and to use them as probes of characteristics which are difficult to measure or strongly biased. In particular, the FP could in principle be a powerful tool to measure distances. But a proper calibration is required, and this appears difficult with its different tilts depending on the galaxy populations.  Interpretation of the FP very often assumes some homology which is defined by \\citet{Gargiulo2009} as: ``systems with density, luminosity and kinematics structures equal over the entire early-type sequence and with constant mass-to-light ratios''. \\citet{vanDokkum2003} define homology in such a way that the evolution of the FP is due only to the evolution of $M/L$. These definitions are certainly linked to the assumption that all early-type galaxies are assembled in the same way, like dissipational mergers \\citep{Robertson2006,Hopkins2008}.  Reality is however more complicated \\citep[see e.g.][]{Bender1992, Jorgensen1996, Borriello2003, vanderMarel2007} and early-type galaxies are very probably not all the result of the same formation process. Basically, what is needed is some invariant that allows us to trace a given object or class of objects through changes due to evolution. This invariant has been hoped to be the FP relation with some universality among a given population, universality which is implied by the definition of homology above and characterized essentially by $M/L$. This provides a rough way of simplifying the many variables that can evolve and may hide important sides of galaxy evolution.  Models assuming only one homologous population have apparently failed to fit the FP in its entire extent. Selection criteria were proposed in order to obtain more homogeneous samples and thus define what could be called a ``purer'' fundamental plane. However, these criteria are necessarily arbitrary, subjective and/or model dependent. The difficulty is that many parameters are known to influence the global shape of the FP. Tilt of the FP \\citep{Robertson2006}, warps, dispersion, changes with redshift, dependence of the mass-distribution on mass \\citep{Nigoche-Netro2009} among others, show that even though the FP looks tight in the $\\log\\sigma$, $<\\mu_e>$ and $\\log{R_e}$ space, additional parameters could still play a role.  The universality of the FP is also questioned: is it a bent plane \\citep{Gargiulo2009} or a bent surface approximated by a collection of planes \\citep{DOnofrio2008}? Obviously, this problem is related to the choice of the sample, that is to the definition of homology. Theorists cope with many parameters that may influence the evolution of a given galaxy, especially when mergers are considered \\citep{Robertson2006}. Testing parameters one after the other, both theoretically and observationally, takes a lot of time, and might partly explain why after so many years, this tight correlation keeps most of its mysteries.  We think that it is time to explore new methodologies to better characterize and understand the FP relation. Multivariate clustering approaches are more objective in selecting really multivariate ``homologous'' sub-populations of galaxies. This requires to explicitly assume that the ``fundamental plane'' is a priori not universal, and to understand it as a correlation in the $(\\log{R_e}, \\log\\sigma,  <\\mu_e>)$ space that could depend on the sub-population. Since galaxies are evolutive objects, homology can be more rigorously defined by 'similarity due to same class of progenitor'. It is the astrophysical equivalent of 'similarity by common ancestry' in cladistics, a statistical method designed to relate evolutionary objects and developed mainly by biologists. The use of many parameters is necessary to find the true homology, and to prevent analogy (same characteristics due to convergent evolution) to yield false lineages of galaxies. Cladistics does not \\emph{assume a priori} properties linked to homology, but rather relies on all the pertinent parameters to \\textit{construct} homologous groups.  In the present paper, we have performed multivariate classifications with two independent approaches, cluster analysis and cladistic analysis, of a carefully chosen sample from the literature. For this first kind of study, we consider only the three parameters of the FP, $\\log\\sigma$, $<\\mu_e>$, $\\log{R_e}$, plus $Mg_2$ that are all given for this homogeneous sample. The first justification is that the number of parameters is small, so it is not necessary to use PCA because we can physically discuss the variation in the light of these few parameters easily. The second advantage is that we are interested to study the evolution of early-type galaxies in the light of their FP properties, and most of the authors have discussed evolution of galaxies on the basis of global FP.  The first approach we have used is by a multivariate technique known as K-means Cluster Analysis using the parameters above. The second approach, known as astrocladistics, is based on the evolutionary nature of both galaxies and their properties \\citep{FCD06,jc1, jc2,DFB09,FDC09}. The clustering technique compares objects for their global similarities, while astrocladistics gathers objects according to their ``histories''. The two techniques are indeed complementary, the first one identifying coherent groups, the second one establishing an evolutionary scenario among groups of objects. They are also totally independent, so that the comparison of their results is extremely instructive from a statistical point of view.   This paper is organized as follows. The data and the different multivariate analyses we performed are described in Sect.~\\ref{multivanal}. Comparison of the structures resulting from these analyses, characterization of the fundamental planes of individual groups, evolution properties within the global FP and properties of the groups, are presented in Sect.~\\ref{results}. The discussion on the meaning of the groups we have found is given in Sect.~\\ref{discussion}, before our conclusion in Sect.~\\ref{conclusion}. We present detailed descriptions of the cluster and cladistic analyses respectively in Appendix~\\ref{appendClus} and Appendix~\\ref{appendClad}, and  provide additional diagrams in Appendix~\\ref{otherfigures}. ",
        "conclusions": "\\label{conclusion}  In this paper, we have reconsidered the study of the so-called fundamental plane of early-type galaxies. We have used two different multivariate clustering tools, cluster analysis and cladistic analysis, to explore the 4-parameter space $\\log\\sigma$, $<\\mu_e>$, $\\log{R_e}$ and $Mg_2$. With both methods, we used both the three first observables, those of the fundamental plane correlation, and the four altogether.  The sample used in our analysis, taken from \\citet{hudson2001}, has 699 objects spread into 56 galaxy clusters. In all the analyses, 4 to 7 groups are found,  which are consistent from both statistical and physical arguments. The very good agreement between our different analyses provides a good confidence in the existence of structures \\textit{within} the FP. We emphasize that no a priori criterion is used to select groups of objects, even in the multivariate space.  In this paper, we focus on the 7 groups (labelled C1 to C7) defined by the cladistic analysis because it additionally provides a diversification scenario linking the groups. Note that since we are in a multivariate space, the wording ``diversification'' is more appropriate than ``evolution'' that applies to a single parameter or that has the general meaning of ``transformation with time''.  The groups define separate regions on the global fundamental plane, not across its thickness. In fact, each group shows its own fundamental plane, which is more loosely defined for less diversified groups. We conclude that the global FP is not a bent surface, but made of a collection of several groups characterizing several fundamental planes with different thicknesses and orientations in the parameter space. In addition, since all groups are present in all galaxy clusters, we conclude that our classification, not the FPs, is universal within the redshift range of the sample (0.007 -- 0.053).   By design in the cladistic analysis, each group supposedly gathers objects sharing a similar history. They are also related by evolutionary relationships that represent a diversification scenario. By rooting the scheme with the group of least metallicity, we find that the two most diversified groups (C6, C7) have the thinnest FP, together with an intermediate one (C3). It probably indicates that the level of diversity is linked to the number and the nature of transformation events like collapse, accretion, interaction and merging, dry and wet, and that the fundamental plane is the result of several such transforming events.  Our groups have distinct multivariate properties that can thus be interpreted in the light of our current knowledge and understanding of galaxy evolutionary processes. Three groups (C1, C2 and C4) probably did not form by major mergers and have been strongly affected by interaction, some of the gas in C2 objects having possibly been swept out. Three other groups, C3, C6 and C7 have been formed by dissipative (wet) mergers because they follow a tight FP relation. In C3, this(ese)) merger(s) must have been quite ancient because of the relatively low metallicity of its galaxies. But contrarily to C7, both C3 and C6 have subsequently undergone dry mergers to increase their masses, more violent in C6 than in C3. Also, in the k-space, the C7 group clearly occupies the region where bulges (of lenticular or spiral galaxies) lie. It could have formed through minor mergers and accretions.   It has been recognized that the properties of the FP depends much on the sample. Our approach strongly associates the ``sample'' with the ``evolutionary group''. Since galaxies are evolutive objects, homology, a concept often used for FP studies, can be more rigorously defined by 'similarity due to same class of progenitor'. Since galaxy histories are so complex, only multivariate studies can objectively construct homologous groups. Cladistics is indeed designed to build homologous groups and provides an evolutionary scenario that relates them. It must be understood that the assembly and interaction histories of nearby galaxies are necessarily complex and comprises often several significant transforming events \\citep[e.g.]{jc2}. The interpretation, based on specific assembly histories of galaxies,  of our seven groups shows that they are truly homologous. Our work also shows that multivariate cluster analysis is able to find homologous groups even though it cannot predict the evolutionary relationships.  The properties of the 7 groups clearly reveal that they differ in assembly histories. Since they have been obtained directly from several observables, the interpretation of the result does not depend on any a priori classification. In particular, it partly matches refined morphological classification because we used 3 parameters (among 4) that are linked to structural properties of galaxies. However our classification is more easily compared to models and numerical simulations. Our work can be readily repeated with additional observables.  In addition, the diversification scenario relating these groups does not depend on models or numerical simulations. The astrocladistic analysis was based on the assumption that the four parameters $\\log{R_e}$, $<\\mu_e>$, $\\log\\sigma$ and $Mg_2$ are evolutive characters, that evolve and can characterize states of evolution of galaxies. The astrophysical interpretation of the diversification scenario and of the evolutionary groups demonstrates a posteriori that this assumption is correct. This is another proof that cladistics can be applied in astrophysics. Since both cluster and cladistic analyses are only valid for the sample in study and the variables used, the present study will be extended to other galaxy samples with more parameters."
    },
    "1005/1005.2253_arXiv.txt": {
        "abstract": "{ We report on the environmental dependence of galaxy properties at $z=2.15$. We construct multi-band photometric data sets in the (proto-)cluster PKS1138-26 field and in the GOODS field.  We then fit spectral energy distributions of the galaxies with model templates generated with the latest stellar population synthesis code and derive physical properties of galaxies from the fits. To quantify the environmental dependence of galaxy properties, a special care is taken of systematic errors -- we use data sets that have almost the same wavelength samplings, use the same code to fit SEDs with the same set of templates, and compare {\\it relative} differences between the two samples. We find that the PKS1138 galaxies have similar ages, shorter star formation time scales, lower star formation rates, and weaker dust extinction compared to the GOODS galaxies at $z\\sim2$. This trend is similar to that observed locally, suggesting that the environmental dependence of galaxy properties is already partly in place as early as $z=2.15$. We show that the PKS1138 galaxies assemble the bulk of their masses $\\sim1$ Gyr earlier than field galaxies, i.e., the galaxy formation depends on environment. Galaxy mergers should frequently occur during the first collapse of clusters and they might play an important role in driving the observed environmental dependence of galaxy properties at $z=2.15$. }{}{}{}{} ",
        "introduction": "The formation and evolution of galaxies in the Universe are dependent on environment in which galaxies live. In the local Universe, red early-type galaxies are the dominant population in rich galaxy clusters, while blue late-type galaxies are the dominant population in the low-density field. Not only galaxy properties, but the formation epoch of galaxies also depends on environment in the sense that cluster galaxies form earlier than field galaxies (e.g., \\citealt{kuntschner02,gebhardt03,thomas05}). Earlier studies concentrated on nearby galaxies, but with the recent advent of large telescopes, environment studies at $z\\sim1$ became possible. Interestingly, the environmental dependence of galaxy properties observed at $z\\sim1$ is already strong; clusters at $z\\sim1$ are dominated by red early-type galaxies (e.g., \\citealt{blakeslee03,nakata05,postman05,lidman08,mei09}, but see also \\citealt{cucciati06}). Also, the formation epoch of cluster galaxies measured at $z\\sim1$ is consistent with that observed locally \\citep{gobat08} . Although there is a clear sign of galaxy evolution between $z=1$ and 0 (e.g., \\citealt{elbaz07,cooper08}), one has to observe galaxies at even higher redshifts to fully quantify the environmental dependence of galaxy formation and evolution.    Because of observational difficulties, only a few high redshift clusters are known so far, the highest redshift cluster being at $z=1.45$ \\citep{stanford06}. Higher redshift galaxies appear fainter and their rest-frame optical light migrates to the near-IR, where the sky background is brighter and it is challenging to observe faint galaxies. Furthermore, high redshift clusters are poor clusters as they are still fast growing according to the dark matter halo growth models \\citep{press74,springel05}. Such poor clusters are difficult to locate due to their weak density contrasts to the general field. There are a number of ways to find high redshift clusters, but one of the proven techniques is to look around high redshift radio galaxies \\citep{miley08}. While not all the radio galaxies are in over-density regions, many of them host clear over-densities of galaxies around them \\citep{venemans07} and they are called proto-clusters. Although over-densities of red massive galaxies are not necessarily confirmed around them, they likely virialize and evolve to clusters at lower redshifts. Among the several proto-clusters reported so far, PKS1138-26 at $z=2.15$ is one of the most promising proto-clusters for its clear over-density of spectroscopically confirmed galaxies by previous studies \\citep{miley08}.    Early studies of the the PKS1138-26 radio galaxy at $z=2.15$ were performed by \\citet{pentericci97,pentericci98}, who reported a clumpy morphology of the radio galaxy. Followed by these initial observations, \\citet{kurk00} first reported an over-density of star forming galaxies around the radio galaxy. This region was then followed up by several authors. \\citet{pentericci00} performed spectroscopic follow-up observations of Lyman $\\alpha$ emitters reported in \\citet{kurk00} and confirmed 14 galaxies close to the radio galaxy redshift. \\citet{kurk04a,kurk04b} carried out further imaging and spectroscopic observations of the field targeting H$\\alpha$ emitters and confirmed another 10 galaxies at the cluster redshift. An X-ray observation has also been performed \\citep{pentericci02} and follow-up spectroscopy confirmed  at least 5 X-ray sources at the cluster redshift \\citep{croft05}. By now, there are more than 20 objects confirmed at the cluster redshift. Detailed analyses of the radio galaxy and the surrounding region with the superb resolution imaging with HST have also been performed \\citep{miley06,zirm08,hatch08,hatch09}, which added further lines of evidence for the forming (proto-)cluster. In fact, \\citet{zirm08} reported on the forming cluster red sequence. Recently, \\citet{doherty09} carried out a near-IR spectroscopic follow-up observation and confirmed two massive red galaxies at the cluster redshift.   Given the the convincing over-density of galaxies and wealth of imaging and spectroscopic data available in the field, PKS1138 is an ideal sample to study the environmental dependence of galaxy properties at this high redshift. In this paper, we perform an extensive analysis of galaxies around PKS1138 to quantify the environmental dependence of galaxy evolution and formation at $z=2.15$.   The layout of the paper is as follows. We summarize our data in Section 2. We then  describe details of our method of fitting spectral energy distributions of galaxies in Section 3. Before presenting our results, we perform extensive sanity checks in Section 4. Section 5 presents physical parameters of galaxies obtained from the fits as a function of environment at $z\\sim2$ and Section 6 discusses implications of our results for galaxy formation. Finally, we summarize the paper in Section 7.   Unless otherwise stated, we adopt H$_0=70\\rm km\\ s^{-1}\\ Mpc^{-1}$, $\\Omega_{\\rm M}=0.3$, and $\\Omega_\\Lambda =0.7$. Magnitudes are on the AB system. We use the following abbreviations: IMF for initial mass function, SED for spectral energy distribution and SFR for star formation rate. ",
        "conclusions": "From the extensive SED fits of galaxies in PKS1138 and GOODS, we have found that massive galaxies in PKS1138 tend to have  \\begin{enumerate} \\item similar ages \\item shorter star formation time scales \\item lower star formation rates \\item lower amounts of dust \\end{enumerate}  \\noindent compared to those in GOODS. The combination of first and second points is interesting. Galaxies in PKS1138 and GOODS start forming stars at a similar epoch in a statistical sense, but PKS1138 galaxies form more rapidly. We recall that our definition of age is the time since the onset of star formation. The third and fourth points are basically by-products of the first and second points. At $z=2$, PKS1138 galaxies have already undergone intense star formation and their SFRs are rapidly declining, while GOODS galaxies are still actively forming stars due to their longer star formation time scales. This results in the lower SFRs in PKS1138 (third point). The fourth point can then be easily understood given the correlation between between SFR and dust amount (lower SFR galaxies have less dust, e.g., \\citealt{hopkins03}).   Cluster galaxies have a shorter star formation time scale -- this is the same trend as observed in $z\\lesssim1$ clusters. For example, \\citet{gobat08} showed the same trend in a $z=1.2$ cluster based on a photo-spectroscopic analysis. It is striking that the trend holds even at $z=2.15$. The difference in the star formation time scale suggests that cluster and field galaxies may form in different ways. Let us discuss the formation and evolution of cluster and field galaxies in detail. We first introduce two parameters to further quantify the galaxy formation. We then extend the discussion and address the origin of the environmental dependence of galaxy properties observed at lower redshifts. We note a caveat here that our results are based on one (proto-)cluster only and the trends we observe may not represent global trends at $z=2$. A larger sample of $z=2$ clusters should be investigated to draw a global picture.   \\subsection{The environmental dependence of galaxy formation}  We have age, star formation time scale, and SFR for each galaxy. Assuming the exponentially decaying star formation histories, we can estimate the SFR at the onset of star formation, which we call initial SFR ($\\rm SFR_0$; see Fig. \\ref{fig:tau_model}). Fig. \\ref{fig:sfr0_mstar} plots  $\\rm SFR_0$ as a function of stellar mass. The cluster galaxies tend to have higher $\\rm SFR_0$ -- the fractions of $\\rm SFR_0>1000M_\\odot\\ yr^{-1}$ are $0.61\\pm0.21$, and $0.39\\pm0.15$ in PKS1138 and GOODS, respectively. The Mann-Whitney test supports this difference (the median probability is $2$\\% and we obtain $<5$\\% probability in 74\\% of the realizations). PKS1138 galaxies have experienced more intense formation histories. Of course, all this is based on the assumption of the exponentially declining SFRs, and we have ignored effects of galaxy-galaxy mergers, which we will discuss later.   If galaxies continue to form stars following the exponential decay, we can derive the time when the galaxies form bulk of their stars. We plot in Fig. \\ref{fig:m80_mstar} the epoch when galaxies form 80\\% of the stars that they would have at $z=0$ ($t_{0.8}$; see Fig. \\ref{fig:tau_model}). PKS1138 galaxies typically formed around $z\\sim3$ or higher, while GOODS galaxies typically formed below $z\\sim3$. We recall that we obtained $z_f\\sim4$ from the location of the red sequence in PKS1138 in Section 5.1, which is in agreement with what we find here. We find that $0.61\\pm0.21$ of galaxies in PKS1138 form the bulk of their stars at $z>3$, while the fraction is $0.17\\pm0.09$ in GOODS. The Mann-Whitney test suggests that they are likely different -- the median probability is $\\sim8\\%$ and we obtained $<5$\\% probabilities in 38\\% of the realization. Interestingly, the ages of the cluster and field galaxies are not very different (Fig. \\ref{fig:physical_params}a). The difference in $t_{0.8}$ is therefore due to their shorter formation time scales. The formation of cluster galaxies is a more intense event, they form in a shorter time scale, and they assemble the bulk of their stars earlier by $\\sim1-2$ Gyr than the field galaxies.  \\begin{figure} \\centering \\includegraphics[width=8cm]{sfr0_mstar.eps} \\caption{ $SFR_0$ is plotted against stellar mass.  The meanings of the symbols are the same as in Fig. \\ref{fig:physical_params}. } \\label{fig:sfr0_mstar} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=8cm]{m80_mstar.eps} \\caption{ $t_{0.8}$ is plotted against stellar mass.  The meanings of the symbols are the same as in Fig. \\ref{fig:physical_params}. } \\label{fig:m80_mstar} \\end{figure}   \\begin{figure} \\centering \\includegraphics[width=8cm]{tau_model3.eps} \\caption{ The median star formation histories of the PKS1138 and GOODS galaxies. The shaded area shows 68 percentile intervals. } \\label{fig:tau_model2} \\end{figure}   To summarize, we measure the median $\\tau$ and SFR for $>10^{11}\\rm M_\\odot$ galaxies and illustrate the differences in star formation histories of galaxies in PKS1138 and GOODS in Fig. \\ref{fig:tau_model2}. This will the summary plot of the paper. As discussed above, PKS1138 galaxies have higher $\\rm SFR_0$ and shorter $\\tau$. PKS1138 experience much more intense galaxy formation at early times than GOODS galaxies.  The galaxy assemblies are completed in a short time, and by the time we observe them ($\\sim2$ Gyr after the onset of star formation as shown in Fig. \\ref{fig:physical_params}a) their SFRs become low. The plot summarizes the points listed at the beginning of this section (except for the fourth one, which is a result of SFR-dust correlation).   The plot indicates that PKS1138 galaxies and GOODS galaxies did not form in the same way.  The formation of cluster galaxies is a more intense event. But, how can massive galaxies form on such a short time scale? In general, star formation activities have negative feedback to themselves. Hot, young stars ionize surrounding gas, which prevents further star formation. Also, supernova explosions of massive stars give negative feedback. If galaxies undergo an AGN phase, it might also give negative feedback. The formation of massive galaxies on a very short time scale may be difficult to understand in this respect, and that makes us speculate that these massive galaxies might have taken another route -- mergers.   Galaxy formation takes place in high peaks of density fluctuations in the Universe. Clusters that we observe today were regions of large-scale over-densities, where many density peaks were embedded. On the other hand, field galaxies form in more isolated peaks. The current galaxy formation theories predict that small galaxies form first and they form progressively massive galaxies via mergers. In the early Universe, cluster galaxies must have had a lot better chances to merge with other galaxies than field galaxies simply because there are more galaxies around them. This accelerated frequency of mergers might be the cause of the rapid formation of cluster galaxies.   A simple scenario of an equal mass merger can be helpful here to show that mergers tend to make $\\rm SFR_0$ higher and push $t_{0.8}$ to higher redshifts. Assuming that two equal mass galaxies merge into a single galaxy, the merger event will double the stellar mass of the galaxy. If the galaxy experiences starburts triggered by the merger, its stellar mass will be more than double and its SFR gets lower after the burst. We then observe this massive galaxy with low SFR, and depending on the time elapsed after the burst, we would fit a short $\\tau$ to this galaxy in order to reproduce its low SFR and large stellar mass. In fact, we find that the secondary burst discussed in Section 4.4 reduces the median $SFR_0$ of the PKS1138 galaxis by $\\sim30\\%$, although it is still high, $SFR_0=1600\\rm M_\\odot\\ yr^{-1}$, which might indicate that galaxies experienced more than one merger. On the other hand, the $SFR_0$ of the GOODS galaxies remain the same within $10\\%$.    We are not yet sure if PKS1138 is a collapsed cluster or currently collapsing proto-cluster.  But, even if it is a cluster, it must be a young system, and it is unlikely that galaxies had enough time to be affected by 'nurture' effects such as ram-pressure stripping. We may suggest that mergers in the early times of the cluster formation may be one of the key processes to establish the environmental dependence of galaxy properties observed at lower redshifts. We will further pursue this point in the following section.   \\subsection{The Nature effects in the environmental dependence of galaxy properties}  Galaxy properties such as colors and morphology are known to depend on environment in which galaxies reside. This environmental dependence is shaped by two effects -- nature and nurture effects. That is, both how galaxies form and how they evolve are important.   Fig. \\ref{fig:m80_mstar} shows that only $\\sim1$ Gyr has passed since the PKS1138 galaxies formed bulk of their stars to the observed epoch. Assuming that PKS1138 is a virialized system, we apply a $2\\sigma$-clipped gapper method \\citep{beers90} and obtain a velocity dispersion of 400 $\\rm km\\ s^{-1}$ using spectroscopic redshifts within 2 arcmin, corresponding to a physical scale of 1 Mpc, from the literature \\citep{pentericci00,kurk04a,croft05}. The virial radius ($r_{200}$) of the cluster is 0.32 Mpc, giving a crossing time scale of $\\sim 1$ Gyr. Therefore, the time elapsed since the bulk formation of the PKS1138 galaxies is comparable to the crossing time scale of this cluster. This suggests that nurture effects may not have had enough time to fully work. They may have affected a fraction of galaxies, but they probably could not change the average properties of galaxies. We may be witnessing the nature effects in shaping the environmental dependence in PKS1138.   We admit that it is not very straightforward to classify nature and nurture effects at this high redshift. Now, let us introduce two kinds of mergers; early-epoch mergers and late-epoch mergers. We refer to mergers occur during the first collapse of clusters as early-epoch mergers, and those occur afterwards as late-epoch mergers. This is to sort out the two effects, early-epoch mergers being nature effects (initial conditions) and late-epoch mergers being nurture effects (environmental effects). Of course, nurture effects also include ram-pressure stripping, harassment, etc. On the other hand, nature effects do not include these processes driven by the deep potential well or intracluster medium.  In this classification, the environmental dependence of PKS1138 is likely due to early-epoch mergers because $\\sim1$ Gyr time is probably not enough for nurture effects to fully work. Galaxy clusters at $z\\sim1$ do not give useful information about nature effects because nurture effects have had enough time to fully operate ($\\sim4$ Gyr since the formation epoch) and it is not straightforward to disentangle the two effects. Our results suggest that nature effects are strong and they form the basis of the environmental dependence of galaxy properties. A way to probe the significance of the nature effects is to quantify morphology of the PKS1138 galaxies. If early-epoch mergers are an important effect, then we expect to observe more signs of recent interactions in PKS1138 than in GOODS (see below). Also, we expect to observe post-starburst galaxies if galaxies undergone interaction triggered starbursts. A near-IR spectroscopic follow-up campaign of PKS1138 is currently underway. We will be able to study spectral properties of the PKS1138 galaxies, which allow us to look deeper into star formation histories.  To sum up, the strong nature effects may shape the environmental dependence of galaxy properties. An ultimate goal of environment studies will be to quantify the relative contribution of nature and nurture effects. But, that will require statistical work on $z\\sim2$ (proto-)clusters. Although the possible significance of nature effects we suggest here is based only on one (proto-)cluster PKS1138, let us further discuss it. It has an interesting implication for the build-up of the cluster red sequence.     \\subsection{The massive end of the cluster red sequence}  The sequence of red early-type galaxies is a ubiquitous feature of galaxy clusters. Over the last few years, there is an accumulating amount of evidence in the literature that the cluster red sequence grows from the massive end to the low-mass end (e.g., \\citealt{tanaka05,tanaka07a,koyama07,tanaka08,gilbank08}). The massive end of the red sequence cannot be formed via a simple fading of blue, star forming galaxies because such massive blue galaxies do not exist even at $z\\sim2$ (see Fig. \\ref{fig:cmd} and discussions in \\citealt{faber07}). We need mergers to form it.   An interesting point here is that we observe the brightest tip of the red sequence in PKS1138, which was also noted by \\citet{zirm08}. The massive end of the red sequence in a young system -- this might be due to early-epoch mergers. Early-epoch mergers might have formed very massive galaxies during the first gravitational collapse of clusters. Early-epoch mergers should occur more frequently in cluster environments than in the field, and that helps explain why we do not observe red sequence in GOODS. As suggested by \\citet{zirm08}, the red sequence in PKS1138 may be being formed or just formed at the time of observation. This view is further supported by our observation that roughly half of the red sequence galaxies have high SFRs (Fig. \\ref{fig:physical_params}). This formation redshift of $z\\sim2$ is in line with predictions from the build-up of the red sequence observed in lower redshift clusters \\citep{tanaka07a}.   We present in Fig. \\ref{fig:acs_blowup} ACS $I$-band images of the bright red galaxies in PKS1138 ($K_s<21$ and $J-Ks>1$). Here we only briefly discuss their morphologies and a detailed study will be presented elsewhere (Zirm et al. in prep). Interestingly, a half of the galaxies show disturbed morphology and/or have nearby companions, lending a support to the picture of the accelerated mergers in clusters. Two out of three apparently disturbed galaxies are detected in MIPS. Early-epoch mergers during the first collapse of clusters may form the brightest end of the red sequence, and at the same time, form the basis of the environmental dependence of galaxy properties.   \\begin{figure*}[tbh] \\centering\\noindent \\includegraphics[width=4.5cm]{acs395.eps} \\includegraphics[width=4.5cm]{acs451.eps} \\includegraphics[width=4.5cm]{acs487.eps}\\\\ \\includegraphics[width=4.5cm]{acs495.eps} \\includegraphics[width=4.5cm]{acs589.eps} \\includegraphics[width=4.5cm]{acs632.eps} \\caption{ ACS $I$-band images of the bright red galaxies. } \\label{fig:acs_blowup} \\end{figure*}    The low-mass end of the red sequence is built up at later times. Nurture effects may come in there. Once clusters form, intracluster gas and deep potential field, in addition to late-epoch mergers, can affect galaxies and terminate their star formation activities. In this way, moderate-low mass galaxies could become red and form the low-mass end of the red sequence. The combination of nature and nurture effects may have conspired to produce the down-sizing behavior of the observed build-up of the red sequence (e.g., \\citealt{tanaka05}).      \\subsection{Proto-clusters at higher redshifts}  Finally, we finish the discussion with a forecast for future (proto-)cluster studies. We can go back in time and see how the PKS1138 field looked like at higher redshifts as we have star formation histories of individual galaxies. A caveat of course is that we ignore all the early-epoch mergers occurred before the time of the observation (the only way to recover that information is to resolve the galaxies into individual stars). Another caveat is that we cannot track spatial positions of galaxies back in time.  But, PKS1138 is likely a collapsing/collapsed system and it may well have been an over-density region already at $z=4$.    Fig. \\ref{fig:lss_evol} shows distribution of galaxies and their star formation rates evolved back to $z=4$. Compared to Fig. \\ref{fig:lss}, galaxies are more actively forming stars on average.  Interestingly, galaxies in PKS1138 are more actively forming stars and there are more starbursting galaxies with $\\rm SFR>100\\ M_\\odot\\ yr^{-1}$ than in GOODS. This is in stark contrast to Fig. \\ref{fig:lss}, where we saw that PKS1138 galaxies have lower SFRs than those in GOODS. We still observe a hint of a galaxy over density in PKS1138 at $z=4$. The plots suggest that, as we approach the formation epoch of clusters (i.e., early phase of the gravitational collapse to a massive cluster halo), we expect to observe an over density of starbursting galaxies. Along with starbursting galaxies, low SFR galaxies already appear in PKS1138, while such galaxies are extremely rare in GOODS. At this redshift, low SFR galaxies and starbursting galaxies may co-exist in a forming cluster.   It is not easy to show how early-epoch mergers change the picture we see here.  But, we still expect to observe starbursting galaxies triggered by early-epoch mergers in collapsing clusters with a higher over-densities of lower-mass, pre-merger galaxies. We deem that observations with existing/future sub-millimeter arrays would be able to discover many forming clusters at high redshifts. In fact, some submm observations of distant radio galaxies have found possible over-densities of dusty starburst populations around them (e.g., \\citealt{debreuck04,greve07}). Full wavelength observations will be essential for future proto-cluster studies at very high redshifts.    \\begin{figure*} \\centering \\includegraphics[height=10cm]{lss_evol_1138.eps} \\includegraphics[height=10cm]{lss_evol_goods.eps} \\caption{ Same as Fig. \\ref{fig:lss}, but galaxies are evolved back in time to $z=4$ assuming the star formation histories obtained from the SED fits. The sizes of the symbols correlate with stellar mass. The top and right axes show comoving scales at $z=4$. } \\label{fig:lss_evol} \\end{figure*}"
    },
    "1005/1005.5229_arXiv.txt": {
        "abstract": "The Lorentz factor (LF) of gamma-ray burst (GRB) ejecta may be constrained by observations of high-energy (HE) spectral attenuation. The recent Fermi-LAT observations of prompt GeV emission from several bright GRBs have leaded to conclusions of unexpectedly large LFs, $\\Gamma>10^3$. Here we revisit this problem with two main concerns. (1) With one-zone assumption where all photons are assumed to be generated in the same region (radius) and time, we {\\em self-consistently} calculate the $\\gamma\\gamma$ optical depth by adopting a target photon spectrum with HE cutoff. We find that this might be important when the GRB LF is below a few hundreds. (2) Recent Fermi-LAT observations suggest that the bulk MeV-range and HE ($\\ga100$~MeV) emission may arise from different regions. We then consider a two-zone case where HE emission is generated in much larger radii than that of the MeV-range emission. We find that the HE emission may be mainly attenuated by MeV-range emission and that the attenuated HE spectrum does not show an exponential spectral cutoff but a slight steepening. This suggests that there may be no abrupt cutoff due to $\\gamma\\gamma$ attenuation if relaxing the one-zone assumption. By studying the spectra of three bright Fermi-LAT GRBs 080916C, 090510 and 090902B, we show that a bulk LF of\\textbf{ $\\Gamma\\sim600$ }can be consistent with observations in the two-zone case. Even lower LFs can be obtained in the multi-zone case. ",
        "introduction": "Relativistic expansion is a key property of gamma-ray bursts (GRBs), and has been confirmed by measurements of radio afterglow sizes, for examples, the indirect estimation by radio scintillation in GRB 970508 \\citep{wkf98} and direct imaging of nearby GRB 030329 \\citep{taylor04}. These observations revealed mildly relativistic GRB ejecta, $\\Gamma\\sim$a few, in the radio afterglow phase. However, it is well believed that GRB ejecta are ultra-relativistic in the beginning-- this is required to solve the so-called \"compactness problem\" \\citep[e.g.,][]{piran99}. The compact GRB source, suggested by the rapid variabilities in MeV light curves, and the huge luminosity suggest hot, optically thick GRB sources, which is in confliction with the nonthermal and hard GRB spectra. Relativistic expansion of the emission region is introduced to solve this problem. In order for the $\\sim100$~MeV photons, as detected by EGRET in several GRBs, to escape from the emission region, avoiding $\\gamma\\gamma$ attenuation, the bulk Lorentz factor (LF) of the emission region is required to be extremely large, $\\Gamma\\ga10^2$ \\citep[e.g.,][]{ls01,krolik91,fenimore93,wl95,bh97}. Recently, the powerful Fermi satellite reveals in much more detail the high-energy (HE) emission from GRBs. Several bright GRBs are reported to show time-integrated spectra extending up to GeV or even tens GeV, without any signs of spectral cutoff. Assuming the $\\gamma\\gamma$ optical depth for these HE photons are below unity, these observations have leaded to even larger bulk LFs, $\\Gamma>10^3$ \\citep{a09a,a09b,a09c}. This is putting the theoretical problem of relativistic jet formation to extremes.  In the previous constraints two assumptions are usually taken. First, all photons, from low to high energy, are produced in the same region and the same time. This \"one-zone\" assumption is not solid, as Fermi observations actually revealed that: the onset of HE emission is delayed relative to MeV emission \\citep[e.g.,][]{a09a,a09b,a09c}; the HE emission lasts longer than MeV emission \\citep[e.g.,][]{a09a,a09b,a09c}; the bulk emission shifts toward later time as the photon energy increases \\citep{a09a} and the shift is longer than the variability times in MeV light curves, as pointed out by \\citep{li10}; some GRBs obviously show distinct HE components with different temporal behaviors \\citep{a09b,a09c}. All these features may imply that different energy photons are produced in different regions.  In particular, the bulk $>100$~MeV emission in GRB 080916C shows $\\sim1$~s shifting relative to MeV emission, which is much longer than the MeV variability time, $<100$~ms as revealed by INTEGRAL \\citep{greiner09}, strongly implying that $>100$~MeV emission is produced in a region of much larger radii than MeV emission's \\citep{li10}. As pointed out by \\cite{lw08}, within the framework of internal shock model, the internal collisions at small radii, which would produce the prompt MeV emission, are expected to lead to \"residual\" collisions at much larger radii, which would produce low-frequency emission. The electrons accelerated by residual collisions at larger radii inverse-Compton scattering the MeV photons and/or double scattering the low-frequency photons could produce HE emission \\citep{li10,zhao10}. In this case, MeV and HE photons are produced in different regions. In the comoving frame of HE emission region, the MeV photons would be collimated other than isotropic, thus the $\\gamma\\gamma$ absorption is angular dependent.  Second, the target photon spectrum is assumed to be extending to infinity. As pointed out by \\cite{li10}, the calculation of $\\gamma\\gamma$ optical depth taking such a target photon filed is obviously not self-consistent, because the HE spectral end should be cut off due to absorption considered in the calculation.  In this paper, we revisit the problem of GRB LF constraint by modifying the above mentioned two assumptions. We consider in \\S2 a one-zone case where the $\\gamma\\gamma$ optical depth is calculated {\\em self-consistently} by assuming a truncated target spectrum, then we consider in \\S3 a simple two-zone case with anisotropic effect on $\\gamma\\gamma$ optical depth taken into account. In \\S4 we studied the spectra of the three bright Fermi-LAT GRBs and constrain their LFs. \\S5 is discussion and conclusions. In the following we assume the concordance universe model with $(\\Omega_m,\\Omega_\\Lambda)=(0.27,0.73)$ and $H_0=71\\rm km\\,s^{-1}Mpc^{-1}$. ",
        "conclusions": "We have revisited in this work the problem of constraining the GRB LFs by the HE attenuation. Although this problem has been considered by many previous works, two concerns that have been ignored in the previous work have been emphasized here. First, we notice that in the one-zone case in order to self-consistently calculate the $\\gamma\\gamma$ optical depth one needs to consider the target photons with HE spectral cutoff, other than extending to infinity. This concern is important when the LFs are below a few hundreds, or when the luminosity of GRBs are low. Second, we relax the one-zone assumption and consider a simple two-zone case where the beaming of target photons in the emission region should be taken into account. Our results show that in the two-zone case, the $\\gamma\\gamma$ absorption does not lead to an abrupt spectral cutoff but a spectral steepening. If the target photon energy distribution is with a power law with photon index $\\bt$ then the spectral slope is changed by a factor of $\\frac1{\\bt}-1$. This also predicts that there should be no spectral cutoff in the GRB spectra if the prompt emission is not produced in one single region.  It should be noted that there are some attempts by other authors to improve the approximation for the optical depth. \\cite{Baring06} concluded that the pair attenuation signature appears as broken power-law rather than exponential cutoff by considering the skin effect and introducing an attenuation descriptor of $1/(1+\\tau)$ instead of $e^{-\\tau}$. \\cite{Granot08} considered the emission zone as a very thin layer producing impulsive emission. They calculated in detail the opacity evolution during a pulse, and claimed that the attenuation signature can be different from that derived from the simple one-zone approximation. Essentially, these two works still concern one-zone problem, with $\\Delta R\\sim R_{\\rm MeV}$. However in the two-zone problem that we considered here, the HE and MeV emission components are emitted at very different radii, with $\\Delta R\\gg R_{\\rm MeV}$, which leads to much smaller optical depth and hence smaller LF at HE emission region.  Furthermore, we take our new concerns to analyze the spectra of the three bright GRBs 080916C, 090510 and 090902B and found that in the two-zone case a LF of $\\Gamma\\sim600$ can still be consistent with the observed spectra. This relaxes the strict requirement, $\\Gamma>10^3$, in one-zone assumption.  We note that in the present observational situation where only tens to hundreds HE photons detected in one GRB, a slight change of the spectral slope is not easy to be identified. A single power law may still fit the HE spectral tail.  We have considered a simple two-zone case here. However the situation can be more complicated. The central engines of GRBs may naturally create variabilities in a wide range of timescales, e.g., from $\\sim1$~ms to $\\sim10$~s. In the framework of internal shock model, this will lead to kinetic dissipation in a wide range of radii. Even in the single-timescale case, the internal collisions will happen as the ejecta expand until the material is distributed with velocity increasing with radius. In such case we will expect multi-zone other than simple two-zone case. The time-integrated spectrum-- note that the time interval with high enough photon statistic is usually much larger than the variability time-- will be contributed by the multiple regions. We also calculate cases with $\\Gamma<600$, the sum of the flux at the HE end can be comparable to the original flux. This means that the spectra can be consistent with a multi-zone case with the LF $\\Gamma<600$.  The formation and acceleration of relativistic collimated GRB jets are open questions. In the standard \"fireball\" model, the thermal pressure can only accelerate the gas up to a LF $\\Gamma\\la10^3$ \\citep[see, e.g.,][]{piran99,li10}. On the other hand, simulations of magnetic-driven jets \\citep[e.g.,][]{mhdjet} can generate jets with the product of the LF and jet opening angle being $\\Gamma\\theta_j\\approx10-30$, which is consistent with pre-Fermi GRB observations. However for the bright Fermi-LAT GRBs, \\cite{cenko10} constrained the jet opening angles by their afterglows, which, combined with the large LF, $\\Gamma>1000$ from $\\gamma\\gamma$ attenuation argument, suggests much larger values of $\\Gamma\\theta_j$. We stress here that if relaxing the one-zone assumption for GRB multi-band emission, LFs with \"normal\" values, say, $\\Gamma\\la600$, can still be consistent with observations. This relaxes further the theoretic problem of jet acceleration.  Recently, a similar paper, \\cite{Zou10}, considering the same two-zone absorption, is now in preprint. The main difference between two papers is the rest frames for the optical depth calculation, i.e., we consider the comoving frame of the HE emission region while they consider the observer frame. They integrate over the region up to $R_{\\max}$, where the HE photon spatially leaves the MeV front, while our integration corresponding to one dynamical time expansion is equivalent to integration up to $2R$, where the generated HE photon doubles its radius. Because both the number density of the target photons and the angle between the travelling directions of HE photon and the MeV front decrease rapidly with radius, the interaction is strongly dominated by those at small radius, and hence the upper limit of the integration is unimportant-- no matter the upper limit is $R_{\\max}(\\gg R)$ or $2R$ the result is practically the same. Furthermore, they use an \"averaged\" optical depth to constrain the LF, which may not be appropriate since we have shown that no sharp cutoff is expected in the two-zone case. We consider more carefully the spectral profile due to suppression. Finally, they neglect the self absorption in the local region which may contribute significant effect as we show."
    },
    "1005/1005.0310_arXiv.txt": {
        "abstract": "{ Blue straggler, which are stars that appear to be younger than they should be, are an important population of unusual stars in both stellar clusters and the halo field of the Galaxy. Most formation scenarios evoke either stellar collisions or binary stars that transfer mass or merge. } { We investigate high-velocity stars in the Galactic halo and perform a spectral and kinematical analysis to shed light on their nature and origin. Here we report that SDSSJ130005.62+042201.6 (J1300+0422 for short) is an A-type star of unusually large radial velocity (504.6 $\\pm$ 5 \\kms). } { From a quantitative NLTE (and LTE) spectral analysis of medium-resolution optical spectra, the elemental composition is derived. Proper motion measurements combined with a spectroscopic distance estimate allow us to determine its present space velocity. Its kinematical properties are derived by integrating the equation of motion in the Galactic potential. } { We find J1300+0422 to be metal poor ([M/H]=$-1.2$) and exhibit an $\\alpha$-element enrichment ($0.3-0.4$~dex) that is characteristic of the halo population, as confirmed by a kinematical analysis of its 3D space motions, which places it on a highly eccentric retrograde Galactic orbit. } { The mass of J1300+0422 (1.15 $\\pm$ 0.10 M$_\\odot$) is higher than the globular cluster turn-off masses indicating that it is a halo blue straggler star. At a Galactic rest-frame velocity of $\\approx$467~\\kms, the star travels faster than any known blue straggler but is still bound to the Galaxy. } ",
        "introduction": "\\label{sec:intro}  Blue straggler stars (BSS) were first discovered as an unusual subclass of stars in globular clusters \\citep{1953AJ.....58...61S}. They lie on or near the main sequence but are more luminous than the turn-off stars indicating that they are of higher mass than the latter. Apparently, BSS are present in all the Galactic globular clusters \\citep{2003ASPC..296..263P}.  Because the stars in a cluster are believed all to have formed at the same time, whereas the stellar turnoff age decreases with mass, BSs should have evolved off the main sequence to become giants and white dwarfs long ago. It is generally believed that blue stragglers are coeval with the remaining stars in the cluster and originate in normal main-sequence stars that gained mass by means of a recent accretion episode. Most formation scenarios evoke mass transfer in and/or the merger of binary stars or collisions of stars.   It remains unclear whether two types of BSS exist in clusters, which may result from the different suggested mechanisms discussed frequently in the literature. The discovery of two distinct sequences of BSS in the globular cluster M~30 \\citep{2009Natur.462..1128N} indicates that indeed both formation mechanisms are at work.  In open clusters, this does not seem to be the case, because \\citet{2009Natur.462..1132N} indentified a high frequency of binaries among the blue stragglers in the open cluster of NGC~188, most having eccentric orbits with periods of about 1000 days. \\citet{2009ApJ...690.1639B} argue that most if not all BSS in open clusters arise from multiple star systems. This suggests that blue stragglers are formed in both ways in star clusters, with collisions/mergers becoming more common with increasing cluster density.    In the Galactic field, it is more difficult to identify BSS because the stellar age cannot be determined. Field blue stragglers are therefore usually identified as metal-poor main sequence objects that are hotter than globular cluster main-sequence turnoff stars. Nevertheless, many blue metal-poor stars $[Fe/H]<-1$ with main-sequence luminosities have been found to be hotter than the main-sequence turnoff of globular clusters and are therefore considered to be field analogs of the cluster BSS population. These metal-poor stars seem to be so numerous that their specific frequency of appearance relative to regular horizontal branch stars was found to be higher than in globular clusters \\citep{2000AJ....120.1014P}. Since the stellar density is low in the field, binary star evolution and mass transfer is probably the most common path of formation among field BSS. \\cite{2001AJ....122.3419C,2005AJ....129.1886C} pointed out that some BS candidates among metal-poor halo main-sequence stars are not binaries, hence binary star evolution and mass transfer may not be the only path. They argue that the apparently single stars could be the partial remnants of an accreted dwarf satellite galaxy whose star formation continued over a long period of time. These metal-poor main-sequence stars are simply therfore young stars.   A detailed spectral analysis of a blue metal-poor star sample in the field by \\cite{2000AJ....120.1014P} measured a very high binary fraction of at least 67\\%, dominated by long period (wide) binaries. This population appears to be very similar to that of the open cluster NGC~188 \\citep{2009Natur.462..1132N}. Assuming a formation by Roche lobe overflow during the red giant branch stage of the primary, \\cite{2000AJ....120.1014P} identified at least half of the blue metal-poor stars in their survey as blue stragglers. This is supported by more recent and much larger surveys for faint blue stars in the halo. For example, \\cite{2008ApJ...684.1143X} found that blue stragglers account for half of their sample of more than 10\\,000 A-type stars.   We have embarked on a search for so-called hyper-velocity stars \\citep[HVS,][]{2009A&A...507L..37T} based on the work of \\citet{2008ApJ...684.1143X}, who presented radial velocities for a large sample of blue stars. Their sample is a mix of blue horizontal branch (BHB), blue straggler, and main-sequence stars with effective temperatures roughly between 7000 and 10\\,000\\,K according to their colours.  To study the stellar motions in 3D, we need to derive their tangential velocities from proper motions and spectroscopic distances. The measurement of proper motions for faint high-velocity stars is the most challenging part of this project because the stars' distances are large and therefore require highly accurate proper motion measurements.    We focused on the fastest stars in terms of large Galactic rest-frame (GRF) velocities with the aim of determining their nature, distance, and kinematics from detailed quantitative spectral analyses and astrometry. First results have already been reported \\citep{2009A&A...507L..37T,2009_J1539}. We found J0136+2425 to be an A-type main-sequence star travelling at $\\approx$590~\\kms, possibly unbound to the Galaxy \\citep{2009A&A...507L..37T}, which makes it an excellent HVS candidate. More importantly, it was inferred to have originated in the outer Galactic rim, nowhere near the Galactic centre, which would be the favoured place of origin if a supermassive black hole acts as a slingshot as suggested by \\citet[][see also \\cite{2009ApJ...690.1639B}]{1988Natur.331..687H}. Amongst the stars of negative GRF velocity, we discovered that J1539+0239 is a BHB travelling with the largest space velocity of any BHB star known so far, which allowed us to place a lower limit on the mass of the Galactic halo \\citep{2009_J1539}.  Here we report that SDSSJ130005.62+042201.6 (J1300+0422 for short) is a metal-poor blue straggler of Population II on a wide retrograde halo orbit that is only marginally bound to the Galaxy in a standard Galactic potential. ",
        "conclusions": "\\label{sec:conclusion}  We have presented a quantitative spectral analysis of a high-velocity star from the sample of faint blue stars in the halo of \\citet{2008ApJ...684.1143X}. Its radial velocity, proper motion, and spectroscopic distance were derived and a detailed kinematical analysis was performed using the Galactic potential of \\citet{1991RMxAA..22..255A} as well as a potential modified one that assumes a more massive dark matter halo.   The metal-poor A-type star J1300+0422 was identified as a blue straggler of 1.15~M$_\\odot$\\ due to its main-sequence gravity. A detailed NLTE analysis was performed, which we compared to the standard LTE approach. Significant differences were found especially for \\ion{C}{i} and \\ion{O}{i}. With its low metallicity of $[Fe/H]=-1.2$ and characteristic enhancement of $\\alpha$-elements, it would fit perfectly into the sample of \\cite{2000AJ....120.1014P}, apart from the huge space velocity of the star. The kinematic characteristics ($U$, $V$, $e$, $J_{\\rm Z}$) confirm the halo membership of J1300+0422 beyond any doubt. In addition, its trajectory continues far out into the halo. Many blue stragglers were found to be long period binaries \\citep[with periods of several $100$\\ to $1000$~days;][]{2005AJ....129.1886C,2009Natur.462..1132N} with low radial velocity semi-amplitudes ($K\\sim 5-10$~\\kms, \\citealp{2001AJ....122.3419C,2005AJ....129.1886C} ). Whether or not J1300+0422 is such a binary needs to be verified by an extensive radial velocity study.  Our kinematical result is limited by the errors in both the spectroscopic distance and the proper motion. ESA's upcoming astrometry mission GAIA will improve the situation because it will provide a parallax measurement with which to check the spectroscopic distance and improve the proper motion of J1300. GAIA will also have an enormous impact on research into blue stragglers in a more general sense as it will provide astrometry of thousands of halo BSS."
    },
    "1005/1005.0242_arXiv.txt": {
        "abstract": "While usually cosmological initial conditions are assumed to be Gaussian, inflationary theories can predict a certain amount of primordial non-Gaussianity which can have an impact on the statistical properties of the lensing observables. In order to evaluate this effect, we build a large set of realistic maps of different lensing quantities starting from light-cones extracted from large dark-matter only N-body simulations with initial conditions corresponding to different levels of primordial local non-Gaussianity  strength $f_{\\rm NL}$. Considering various statistical quantities (PDF, power spectrum, shear in aperture, skewness and bispectrum) we find that the effect produced by the presence of primordial non-Gaussianity is relatively small, being of the order of few per cent for values of $|f_{\\rm NL}|$ compatible with the present CMB constraints and reaching at most 15-20 per cent for the most extreme cases with $|f_{\\rm NL}|=1000$. We also discuss the degeneracy of this effect with the uncertainties due to the power spectrum normalization $\\sigma_8$, finding that an error in the determination of $\\sigma_8$ of about 3 per cent gives differences comparable with non-Gaussian models having $f_{\\rm NL}=\\pm 1000$. These results suggest that the possible presence of an amount of primordial non-Gaussianity corresponding to $|f_{\\rm NL}|=100$ is not hampering a robust determination of the main cosmological parameters in present and future weak lensing surveys, while a positive detection of deviations from the Gaussian hypothesis is possible only breaking the degeneracy with other cosmological parameters and using data from deep surveys covering a large fraction of the sky. ",
        "introduction": "In recent years, the interest for an accurate measurement of the amount of non-Gaussianity present in the primordial density field has largely increased. The main reason is that this test is now considered not only a general probe of the inflationary paradigm, but also a powerful tool to constrain the plethora of its different variants. Only the most standard slow-rolling models based on a single field produce in fact almost uncorrelated fluctuations, which is the motivation of the common assumption (and large simplification) that their distribution is Gaussian. In general, small deviations from Gaussianity are predicted even for the simplest inflationary models, while non-standard models, like the scenarios based on the curvaton, on the inhomogeneous reheating and on the Dirac-Born-Infeld inflation allow much more significant departures \\citep[see][and references therein]{Bartolo2004}.  It has become common to quantify the level of primordial non-Gaussianity adopting the dimensionless non-linearity parameter $f_{\\rm NL}$ \\citep[see, e.g.,][]{Salopek1990,Gangui1994,Verde2000,Komatsu2001}, that measures the importance of the quadratic term in a sort of Taylor expansion of the gauge-invariant Bardeen potential $\\Phi$\\footnote{We recall that on scales smaller than the Hubble radius $\\Phi$ corresponds to the usual Newtonian peculiar potential (but with changed sign).}: \\begin{equation} \\Phi=\\Phi_{\\rm L}+f_{\\rm NL} (\\Phi^2_{\\rm L}-\\langle \\Phi^2_{\\rm L}\\rangle)\\ ; \\end{equation} here $\\Phi_{\\rm L}$ represents a Gaussian random field. Hereafter we will adopt the so-called large-scale structure (LSS) convention, where $\\Phi$ is linearly extrapolated to the present epoch\\footnote{With the cosmological parameters adopted in this paper, this corresponds to values for $f_{\\rm NL}$ larger by a factor of $\\approx 1.3$ with respect to the so-called CMB convention, where $\\Phi$ is instead extrapolated at $z=\\infty$.}. Moreover we will consider only the so-called local shape for non-Gaussianity, in which the bispectrum signal is larger on squeezed triangle configurations. For more details about other possible shapes we refer to \\cite{Bartolo2004,Verde2009,Bartolo2009.1} and references therein.  At present the most stringent constraints on $f_{\\rm NL}$ come from the cosmic microwave background (CMB) data. Their discriminating power derives from the fact that its temperature fluctuations trace the density perturbations before the gravitational non-linearities modify their original distribution. Whatever is the specific test adopted, all CMB analysis consistently allow only very small deviations from Gaussianity: for instance, analyzing the recent WMAP data, \\cite{Komatsu2010} found that $f_{\\rm NL}$ varies between -13 and 97, while \\cite{Smith2009} found $-5<f_{\\rm NL}<104$ (at 95 per cent confidence level); see also \\cite{Komatsu2009}.  Alternative and complementary constraints on $f_{\\rm NL}$ can in principle be derived analyzing the LSS \\citep[for recent reviews, see, e.g.][]{Verde2010,Desjacques2010}. As already evident from the first generation of non-Gaussian N-body simulations \\citep{Messina1990,Moscardini1991,Weinberg1992}, the presence of a positive (negative) skewness in the PDF of the primordial density field tends to favor (disfavor) the formation of cosmic structures, inducing a different timing in the whole process of gravitational instability. However, to be fully exploited, this approach needs to be complemented with reliable methods to disentangle from the primordial signal the non-Gaussian features introduced by the late non-linear evolution and by the possible presence of a non-linear bias factor. While lensing statistics entirely avoid the latter problem, to attack the former one it is necessary to make use of both analytic techniques \\citep[like high-order perturbation theory and the Time-Renormalization Group approach; see, e.g.,][]{Taruya2008,Bartolo2009.2} and full N-body simulations \\citep[see, e.g.,][]{Kang2007,Grossi2007,Dalal2008,Viel2009,Desjacques2009,Pillepich2009,Grossi2009.1} to properly calibrate the theoretical predictions in the non-linear regime. The large amount of theoretical work recently done in this direction allowed to better understand what is the size of the effects on the abundance of non-linear structures \\citep{Matarrese2000,Verde2000,Mathis2004,Kang2007,Grossi2007,Grossi2009.1,Maggiore2009,Roncarelli2009}, on the halo biasing \\citep{Dalal2008,McDonald2008,Fedeli2009.1,Carbone2010}, on the galaxy bispectrum \\citep{Sefusatti2007.1,Jeong2009,Nishimichi2009.2}, on the mass density distribution \\citep{Grossi2008,Lam2009.1,Lam2009.2}, on the topology \\citep{Matsubara2003.1,Hikage2008}, on the integrated Sachs-Wolfe effect \\citep{Afshordi2008,Carbone2008}, on the Ly-$\\alpha$ flux from low-density intergalactic medium \\citep{Viel2009}, on the 21 cm fluctuations \\citep{Cooray2006,Pillepich2007} and on the reionization process \\citep{Crociani2009}. The first attempts of an application to real observational data gave very encouraging results: \\cite{Slosar2008}, combining the bias measurements for two samples of luminous red galaxies and quasars, found $f_{\\rm NL}=48^{+55}_{-74}$; \\cite{Afshordi2008}, studying the integrated Sachs-Wolfe effect (ISW) in the NVSS survey, derived $f_{\\rm NL}=354\\pm 165$; all error bars are at 2-$\\sigma$ level\\footnote{We report the values as revised by \\cite{Grossi2009.1} to include a correction mimicking the ellipsoidal collapse and converted to the LSS convention.}. Very recently \\cite{Xia2010} found $32 < f_{\\rm NL}< 152$ from the analysis of the auto-correlation of the brightest NVSS sources on angular scales of several degrees.  In this paper we will focus on estimating the weak lensing signals in scenarios with primordial non-Gaussianity. Being based on the measurement of the shear effect produced by the intervening large-scale structure of the universe on the images of background galaxies, gravitational lensing is a direct probe of the total matter distribution. For this reason it is considered one of the most powerful tools to constrain the main cosmological parameters and many dedicated projects are in progress or under study. A very exciting perspective is certainly opened by the ESA Cosmic Vision project EUCLID \\citep{Laureijs2009}, currently under study: the goal of its wide survey is to obtain the shear measurements for about 40 galaxies per arcmin$^2$ on the entire extragalactic sky with galactic latitude $b>30$ (approximately 20000 deg$^2$). The possibility of using the weak lensing signals to constrain also the amount of primordial non-Gaussianity has been already explored by different authors. \\cite{Amara2004} used a generalized halo model to study the impact on the estimates of the power spectrum normalization $\\sigma_8$ of primordial non-Gaussianity, modeled assuming various lognormal distributions for the density field. More recently, \\cite{Fedeli2010} computed the power spectrum of the weak cosmic shear for non-Gaussian models with different values of $f_{\\rm NL}$. In particular, they improved the halo model including more accurate prescriptions for its ingredients (mass functions, bias and halo profile), calibrated on the last generation of non-Gaussian N-body simulations. The application of this model to a survey having the expected characteristics of the EUCLID project showed the possibility of a significant detection of non-Gaussianity at the level of $|f_{\\rm NL}|\\approx$ few tens, once the remaining parameters are held fixed.  In this paper we investigate weak lensing statistics in non-Gaussian scenarios using numerical rather than analytical tools. Specifically, we will create weak lensing maps performing ray-tracing simulations through very deep light-cones extracted from high-resolution N-body simulations. The advantage of this approach is twofold. First of all, N-body simulations permit to fully account for the non-linear evolution which is usually modeled less accurately by analytical means. Second, numerical experiments allow us to extract a large set of realistic weak lensing maps that can be used for better evaluating the statistical robustness of the results.  The main goal of our numerical work is to figure out what are the observational evidences of the presence of some level of primordial non-Gaussianity, as quantified by the $f_{\\rm NL}$ parameter. In particular we will compute a large set of weak lensing statistics in models with different $f_{\\rm NL}$ and we will quantify the deviations from the corresponding results in the Gaussian scenario. This is important not only to address the possibility of a positive detection with future data, but also to establish at which level an amount of primordial non-Gaussianity compatible with the present observational constraints can hamper an accurate measurement of the other cosmological parameters. We must recall that Gaussian initial conditions are virtually always assumed in their practical derivation.  The plan of this paper is as follows. In Section~\\ref{sect:lensing} we review the basis of the lensing formalism necessary to the present work. In Section~\\ref{sect:simulations} we describe the cosmological simulations and the numerical procedure to build the lensing maps. In Section~\\ref{sect:results} we present our main results about the statistical properties of the different lensing quantities investigated: the probability distribution function, the third-order moment, the power spectrum and the bispectrum. In Section~\\ref{sect:sigma8} we compare the effects produced by primordial non-Gaussianity to the uncertainties related to power spectrum normalization $\\sigma_8$. Finally, in Sect.~\\ref{sect:conclusions} we draw our conclusions. ",
        "conclusions": "\\label{sect:conclusions}  In this work, we used the outputs of N-body simulations to create a large set of realistic mock maps for several lensing quantities (deflection angle, effective convergence, shear and the two components of the flexion) in the framework of cosmological models with different amount of primordial non-Gaussianity, quantified using the dimensionless parameter $f_{\\rm NL}$. In particular we considered several statistical properties (PDF, power spectrum, shear in aperture, skewness and bispectrum) and compared the results with the corresponding ones obtained in the Gaussian case. Our main results can be summarized as follows.  \\begin{itemize} \\item For all quantities here considered the effect produced by the presence of primordial non-Gaussianity is relatively small, amounting to differences of few per cent for $|f_{\\rm NL}|=100$, 5-10 per cent for $|f_{\\rm NL}|=500$, and 15-20 per cent for $|f_{\\rm NL}|=1000$. These results are in good agreement with the analytic predictions presented in \\cite{Fedeli2010}. \\item The largest effects are visible on small scales (i.e. for large multipoles $l>1000$), where, however, also non-linearity can produce strong effects which have to be accurately modeled. \\item The most promising statistical tests to search for imprints of primordial non-Gaussianity are the (convergence and shear) power spectra and the (convergence) bispectrum, thanks to the smaller size of their error bars at the relevant scales. \\item The differences of the various PDFs in both rare-event tails are also important, but their discriminating power is reduced by the poor statistics and by the high-level of noise. \\item We compared the effects produced by the primordial non-Gaussianity with the uncertainties due to the power spectrum normalization: an error in the determination of $\\sigma_8$ of about 3 per cent gives differences comparable with the non-Gaussian models with $f_{\\rm NL}=\\pm 1000$, while for more realistic non-Gaussian models with $f_{\\rm NL}=\\pm 100$ the effects is even larger than the one induced by $\\Delta \\sigma_8=0.01$. \\end{itemize}  As said, a significant covariance exists between primordial non-Gaussianity and fundamental cosmological parameters, especially with $\\sigma_8$. The smallness of non-Gaussian effects found in our analysis means that one can obtain a precise estimate of these parameters despite their covariance with $f_{\\rm NL}$. As an example, ignoring the possible presence of a primordial non-Gaussianity with $f_{\\rm NL}=100$, consistent with current observational constraints, would induce a mere 0.2 per cent uncertainty in the estimate of $\\sigma_8$, i.e. well below the current 1-$\\sigma$ error. The same argument, in reverse, is telling us that the search for non-Gaussian features in the weak-lensing statistics is a very challenging task. According to our results, it can be successfully completed only if errorbars can be significantly reduced. For the purpose of discriminating among competing models the reduction needs not to be dramatic. Indeed, as pointed out by \\cite{Fedeli2010}, the fact that deviations from Gaussianity are small but systematics allows one to estimate $f_{\\rm NL}$ by adding the statistical information from different angular bins.  A more significant reduction in the errorbars is required to break the degeneracy between non-Gaussianity with cosmological parameters, say $\\sigma_8$. Indeed, as shown in Fig.~\\ref{fig:sigma8}, removing such degeneracy requires to compare information from a limited number of bins at very different scales.  As shown by \\cite{Fedeli2010}, both tasks could be achieved with next-generation, all-sky surveys. A good example is certainly represented by the ESA mission EUCLID \\citep{Laureijs2009}. According to the analysis made by \\cite{Fedeli2010} and based on analytic predictions for the power spectrum quantitatively confirmed in this numerical work, we expect that the quality and the quantity of the EUCLID data will allow to constrain $f_{\\rm NL}$ at the level of few tens, opening the possibility of discriminating between the various inflationary models."
    },
    "1005/1005.2592_arXiv.txt": {
        "abstract": "{We demonstrate the unique capabilities of {\\it Herschel} to study very young luminous extragalactic young stellar objects (YSOs) by analyzing a central strip of the Large  Magellanic Cloud obtained through the HERITAGE Science Demonstration Program. We combine PACS 100 and 160, and SPIRE 250, 350, and 500  $\\mu$m photometry with 2MASS (1.25-2.17  $\\mu$m) and {\\it Spitzer} IRAC and MIPS  (3.6-70  $\\mu$m) to construct complete spectral energy distributions (SEDs) of compact sources. From these, we identify 207 candidate embedded YSOs in the observed region, $\\sim$40\\% never-before identified. We discuss their position in far-infrared color-magnitude space, comparing with previously studied, spectroscopically confirmed YSOs and maser emission.  All have red colors indicating massive cool envelopes and great youth. We analyze four example YSOs, determining their physical properties by fitting their SEDs with radiative transfer models.  Fitting full SEDs including the {\\it Herschel} data requires us to increase the size and mass of envelopes included in the models.  This implies higher accretion rates ($\\gtrsim$10$^{-4}$M$_\\odot$yr$^{-1}$), in agreement with previous outflow studies of high-mass protostars. Our results show that {\\it Herschel} provides reliable longwave SEDs of large samples of high-mass YSOs; discovers the youngest YSOs whose SEDs peak in {\\it Herschel} bands; and constrains the physical properties and evolutionary stages of YSOs more precisely than was previously possible.} ",
        "introduction": "The proximity of the Magellanic Clouds offers a unique opportunity to analyze the complete inventory of luminous YSOs over an entire galaxy.  With known YSO distances, luminosities, masses, and mass accretion rates can all be well-defined.  Comparison of the properties of YSOs in the  Magellanic Clouds and in the Milky Way can reveal differences in star formation physics due to metallicity and environment.  Using the {\\it Spitzer} SAGE (\\textquotedblleft Surveying the Agents of Galaxy Evolution\\textquotedblright) Survey of the Large Magellanic Cloud \\citep[LMC;][]{meixner06}, Whitney et al. (2008; W08) and Gruendl \\& Chu (2009; GC09) discovered $\\sim$1800 massive YSO candidates in the LMC (a 90-fold increase over previous work). {\\it Spitzer} studies selected sources using colors and SEDs at wavelengths $\\leq$24$\\;\\mu$m (where {\\it Spitzer} can resolve individual YSOs), requiring a detection at 4.5$\\;\\mu$m or shorter in most cases.  These surveys thus missed the youngest, most embedded YSOs that can only be detected at longer wavelengths.  The {\\it Herschel} Space Observatory \\citep{pilbratt10} has the spatial resolution required to study individual sources at $\\lambda\\gtrsim$50 $\\mu$m (from $\\sim$1.3 pc at 70 $\\mu$m to $\\sim$8.7 pc at 500 $\\mu$m for a distance of 50 kpc, \\citealt{schaefer2008}). The least-evolved massive protostars are characterized by cold dust temperatures probed at far-infrared (far-IR) wavelengths, and are expected to be $\\sim$10$^3$ times brighter at 100 $\\mu$m than at 5 $\\mu$m \\citep{whitney04,molinari08}, making {\\it Herschel} extremely effective at detecting those youngest YSOs.  With {\\it Herschel}, we not only discover new objects but also better characterize {\\it Spitzer}-identified YSO candidates.  The {\\it Herschel} data constrain the models of these sources and improve estimates of such physical parameters as total luminosity, stellar mass, and total dust mass. We demonstrate these capabilities by studying a strip across the LMC observed as part of the Science Demonstration Program (SDP) -- the first part of the {\\it Herschel} Key Program \\textquotedblleft HERschel Inventory of the Agents of Galaxy Evolution\\textquotedblright\\ (HERITAGE; \\citealt{meixner10}) in the Magellanic Clouds.  The strip was mapped in the PACS 100 and 160$\\;\\mu$m bands \\citep{poglitsch10} and SPIRE 250, 350, and 500$\\;\\mu$m bands \\citep{griffin10}. ",
        "conclusions": "We show that {\\it Herschel} far-IR photometry is very effective in identifying YSOs in the LMC. Adding {\\it Herschel} data to existing {\\it Spitzer} and near-IR observations results in significantly improved analysis of YSOs, as summarized in these key findings: \\begin{list}{$\\circ$}{\\setlength{\\topsep}{0in}\\setlength{\\parsep}{0in}} \\item Nearly all sources detected by {\\it Herschel} are Stage~0/I, very young, with a high ratio of circumstellar to stellar mass. \\item Previously studied warm sources such as YSO-1 require more circumstellar dust to fit the {\\it Herschel} data, implying a less evolved state than would be inferred from {\\it Spitzer} data alone. \\item {\\it Herschel} photometry significantly constrains our SED fits, decreasing the range of circumstellar dust masses and evolutionary states consistent with the measurements. \\item Many sources require even more cold circumstellar dust than is present in our original model grid, motivating improvements to our models. \\end{list} Our observations prove  that {\\it Herschel} can offer us,  for the first time, an inventory of the earliest stages of protostellar formation throughout an entire galaxy. \\\\"
    },
    "1005/1005.2840_arXiv.txt": {
        "abstract": "{% V391~Peg (HS~2201+2610) is an extreme horizontal branch subdwarf B (sdB) star, it is an hybrid pulsator showing $p$- and $g$-mode oscillations, and hosts a 3.2/sin$i$ M$_{Jup}$ planet at an orbital distance of about 1.7 AU. In order to improve the characterization of the star, we measured the pulsation amplitudes in the {\\it u'g'r'} SLOAN photometric bands using ULTRACAM at the William Herschel 4.2 m telescope and we compared them with theoretical values. The preliminary results presented in this article conclusively show that the two main pulsation periods at 349.5 and 354.1~s are a radial and a dipole mode respectively. This is the first time that the degree index of multiple modes has been uniquely identified for an sdB star as faint as V391~Peg (B=14.4), proving that multicolor photometry is definitely an efficient technique to constrain mode identification, provided that the data have a high enough quality.} ",
        "introduction": "About half of field sdB stars, which reside in binary systems, can form through common envelope ejection or stable Roche lobe overflow (Han et al. 2002, 2003). It is more difficult to explain the formation of a single sdB star. Two scenarios are possible: the merger of two low-mass helium white dwarfs and an early hot helium flash; but both are not fully consistent with the observations. A recent review on these arguments is given by Heber (2009). Another possibility, suggested by Soker (1998), is that the huge mass loss needed to form an sdB star is triggered by low-mass bodies, planets or brown dwarfs (BDs). Although this possibility has not been tested by detailed models yet, a planet to the pulsating sdB star V391~Peg (Silvotti et al. 2007) and three circumbinary planets/BDs to the eclipsing sdB+M systems HW~Vir (Lee et al. 2009) and HS~0705+6700 (Qian et al. 2009) suggest that sdB planets/BDs could be a relatively common phenomenon (see also Bear \\& Soker 2010). A systematic search for sdB substellar objects around 4 sdB stars using the timing method is the main  goal of the EXOTIME project (Schuh et al. 2010, Benatti et al. 2010).  \\vspace{1mm}  V391~Peg (HS~2201+2610) is a particularly interesting system formed by an sdB star and a 3.2/sin$i$ M$_{Jup}$ planet orbiting the host star in 3.2 years at a distance of about 1.7 AU (Silvotti et al. 2007). The sdB star is a hybrid pulsator showing $p$ and $g$-mode oscillations at the same time ({\\O}stensen et al. 2001, Lutz et al. 2009), offering a unique opportunity to characterize the host star through asteroseismic methods. A preliminary mode identification of the higher pulsation frequencies ($p$-modes) was proposed in Silvotti et al. (2002): the two main pulsation periods of 349.5 and 354.1 s could be reproduced with $l$=0 ($k$=1) and $l$=1 ($k$=1) respectively. However this solution was not unique due to the small number of detected frequencies and other solutions could not be excluded.  \\vspace{13.6mm} ",
        "conclusions": "Thanks to the high quality of the data, this is the first time that the mode degree index has been uniquely identified from multicolor photometry for the two main modes of a star as faint as V391~Peg (V=14.6). To our knowledge, conclusive results were obtained only for two brighter stars: KPD~2109+4401 (V=13.4, Randall et al. 2005, see also Jeffery et al. 2004) and Balloon 090100001, the brightest known sdBV star with V=12.1 (Baran et al. 2008, Charpinet et al. 2008).  The results reported in this article confirm that multicolor photometry can set useful identification constraints on the pulsation modes of sdB rapid pulsators, provided that the data have a high enough quality. ULTRACAM on a 4~m class (or larger) telescope is an ideal instrument for such studies."
    },
    "1005/1005.3594_arXiv.txt": {
        "abstract": "{A phenomenological relationship between oscillations in a sunspot and quasi-periodic pulsations in flaring energy releases at an active region above the sunspot, is established. The analysis of the microwave emission recorded by the Nobeyama Radioheliograph at 17~GHz shows a gradual increase in the power of the 3-min oscillation train in the sunspot associated with AR 10756 before flares in this active region. The flaring light curves are found to be bursty with a period of 3 min. Our analysis of the spatial distribution of the 3-min oscillation power implies that the oscillations follow from sunspots along coronal loops towards the flaring site. It is proposed that quasi-periodic pulsations in the flaring energy releases can be triggered by 3-min slow magnetoacoustic waves leaking from sunspots. } ",
        "introduction": "The problem of the transfer of energy, momentum, and information from sub-photospheric solar regions to both the corona and the solar wind is one of the most difficult in solar and stellar physics. Magnetohydrodynamic (MHD) waves are believed to play a key role because the waves are natural carriers of energy, momentum, and information (e.g., Erd{\\'e}lyi 2006 for a recent review). In addition, the guided nature of the wave propagation opens up very interesting perspectives for fixing and tracing the energy transfer channels highlighted by the waves. In general, MHD waves can be guided by inhomogeneities in the characteristic MHD speeds (Alfv\\'en, fast and slow) as well as the magnetic field itself, which are found in all regions of the solar atmosphere.  At the chromospheric level, one of the most pronounced wave phenomena are the 3-min oscillations over sunspots (see, e.g. Bogdan \\& Judge 2006 for a recent review), usually detected as intensity oscillations in visible light, UV, and EUV spectral lines, as well as in microwave band (e.g., Shibasaki 2001, Gelfreikh et al. 1999, Nindos et al. 2002) and in the dm-radio flux records (M\\'esz\\'arosov\\'a et al. 2006). Three-minute oscillations in sunspots are believed to be associated with slow magnetoacoustic waves (e.g., Zhugzhda 2008). Outwardly propagating compressible waves of the same periodicity are also seen in both the EUV 171\\AA\\ and 195\\AA\\ bandpasses in the magnetic fan structures situated over sunspots (e.g., De Moortel 2006 for a review). The projected phase speed of these waves is subsonic, and the waves are seen to propagate along the plasma channels elongated along the coronal magnetic field lines, and hence are interpreted as slow magnetoacoustic waves. Compressible 3-min waves observed at the same location in both EUV bandpasses show a high degree of correlation (King et al. 2003). The relationship between these waves and 3-min oscillations in sunspots remains unclear. The understanding of the propagation of 3-min oscillations through the solar atmosphere is one of the most important problems of solar physics. Its understanding will probably indicate the nature and properties of the plasma channels that transfer these waves into the corona, and hence the connectivity of different layers of the atmosphere. The role played by 3-min oscillations in the corona is also of interest, in particular, the relationship between 3-min oscillations in sunspots and the flaring activity in the active regions (AR) above the sunspots. A possible indication of such a relationship was mentioned in Gelfreikh (2002).  Wave and oscillatory phenomena in various parts of the solar atmosphere can trigger and modulate bursty energy releases, e.g., solar flares. In this case, the periodicity of the oscillations will be evident in the flaring light curves as quasi-periodic pulsations (QPP). This can be achieved by several mechanisms. In the scenario proposed by Nakariakov et al. (2006), energy of transverse (kink or sausage) oscillations of coronal loops can periodically leak to a magnetic neutral point or line situated nearby. The incoming fast magnetoacoustic wave refracts towards the neutral point, experiencing focussing and steepening. This periodically generates very sharp spikes of electric current density in the vicinity of the neutral point, which in turn can be affected by current driven plasma micro-instabilities. The instabilities can cause the onset of micro-turbulence and hence enhance the plasma resistivity by several orders of magnitude. This would lead to periodic triggering of magnetic reconnection and hence the manifestation of the loop oscillations as periodic variation in the flaring light curve.  A compressible wave can periodically trigger magnetic reconnection not only by periodic current density spikes, but also by the variation in the plasma density in the vicinity of the reconnection site. This possibility was modelled numerically by \\cite{2006SoPh..238..313C} in interpreting 3--5~min periodicity detected in repetitive bursts of explosive events in the transition region (Ning et al. 2004). Density variations result in a periodic variation in the electron drift speed. Depending upon the ratio of electron to proton temperatures, the value of the speed controls the onset of the Buneman or ion-acoustic instabilities and hence anomalous resistivity. The periodic onset of the anomalous resistivity triggers periodic energy releases. Transverse compressible waves may also directly trigger magnetic reconnection causing transition region explosive events, by changing the magnetic field strength (Doyle et al. 2006). Longitudinal, e.g., acoustic waves can also modulate flaring energy releases, either directly by the modulation of the drift velocity or the modulation of gyrosynchrotron emission efficiency (Nakariakov \\& Melnikov 2006), or indirectly e.g., by means of centrifugal conversion into fast magnetoacoustic waves on the curved magnetic field lines (Zaitsev \\& Stepanov 1989).  Kislyakov et al. (2006) analysed 15 flares observed in the 37~GHz band with the Mets\\\"ahovi radio telescope (Finland) with the use of the \\lq\\lq sliding window\" Fourier transform and the Wigner--Ville nonlinear transform. The telescope spatial resolution is 2.4~arc min, the sensitivity is about 0.1 sfu, and the time resolution was higher than 0.1~s. During 13 events (about 90\\%), a 5-min periodic modulation of the emission intensity was detected with the frequency of $3.2\\pm 0.37$~mHz. In addition, a shorter period (about 1~s) signal was detected, which was found to be frequency modulated with the same 5-min period. In the development of this study (Za{\\u i}tsev \\& Kislyakov 2006), simultaneous modulation of the microwave emission by three low frequency signals with periods of 3.3, 5, and 10~min was observed in 30\\% of the analysed outbursts. It was suggested that the detected modulation was caused by the parametric resonance between 5-min velocity oscillations in the solar photosphere and natural acoustic oscillations of coronal magnetic loops modulating the microwave emission. The detected periods of 5, 10, and 3 min were interpreted to correspond to the pumping frequency, its subharmonic, and its first upper frequency of parametric resonance, respectively. Confirmation of this finding with a different instrument is required. The physical mechanisms responsible for the appearance of 3 and 5 min flaring QPP cannot be determined without spatial information.  In general, spatial information, such as the spatial size and shape of the region occupied by an oscillation, and the distribution of the oscillation power, phase, and spectrum over the source, is crucial to establishing the nature of the QPP (e.g., Grechnev et al. 2003; Melnikov et al. 2005). Novel imaging data analysis techniques developed in solar physics have been shown to allow one to exploit the full potential of the spatially resolving observations. Grechnev (2003) proposed creating a 2D broadband variance map, representing the overall dynamics of an analysed event, from microwave correlation data cubes obtained with NoRH. This approach allows us to ascertain the spatial locations of faint variable microwave emission sources. Nakariakov \\& King (2007) designed a coronal-periodmapping technique, which reduces 3D imaging data cubes (2D in space and time) to a sequence of static maps inferring collective oscillations of extended (larger than the pixel) coronal structures. This approach was successfully applied by Inglis et al. (2008) to the study of single-periodic pulsations in a large off-limb flaring loop seen in the microwaves with NoRH. However, period-mapping does not provide any phase information, and does not allow detailed studying of multi-periodic or non-stationary phenomena. Complex temporal and spatial features of oscillatory processes in imaging datasets can be studied with the pixelised wavelet filtering (PWF) technique (Sych \\& Nakariakov 2008). This approach produces 4D (2D in space, time and frequency) data cubes providing information about the time modulation of oscillatory signals, their coupling, and their evolution. In particular, this method allows one to obtain information about the spatial structure of narrowband and broadband time signals, as well as the analysis of the signal integrated over the whole field-of-interest. The practical implementation of PWF consists of several steps: \\begin{itemize} \\item The images in the time sequence are coaligned, removing the spatial mismatch between consecutive images. \\item The object of investigation (e.g., the period of the oscillations) and the field-of-interest (FOI) are selected. \\item Construction of the variance map of the FOI, determining the spatial distribution of the integrated power of the time signal. \\item Direct wavelet transform in the time domain, filtering out the spectral components of interest, and inverse wavelet transform. The resultant data cube contains the time variation in the signal in the prescribed spectral band. \\end{itemize} In addition, global wavelet spectra can be calculated for each pixel, resembling the construction of a periodmap. PWF is a convenient tool for establishing a relationship between different spatially-separated oscillatory processes.  Analysis of the possible relationship between flaring energy releases and dynamical processes in the lower regions of the solar atmosphere can shed light on the triggering of the energy releases and hence the basic physical processes responsible for them. The spatially-resolved analysis of this relationship can also infer the atmospheric connectivity channels. The aim of the paper is to establish a phenomenological relationship between dynamical processes occurring in a sunspot and in flares, within the active region (AR) linked to the sunspot, and identify the channels of the connectivity. We demonstrate that 3-min oscillations of a sunspot appear to be present in the microwave emission associated with the flaring activity over the sunspot. ",
        "conclusions": "The microwave light curves of solar flares on 2005 April 28 and 2005 May 4 contain pronounced variations with periods of about 3 min. This behaviour indicates that there is an apparent relationship with 3-min oscillations in the sunspot situated close to the flare sites.  The aim of this paper was to understand this relationship. Our analysis of dynamical features in the microwave, EUV, white light, and X-ray imaging data of AR~10756 acquired during its passage through the solar disk from 2005 April 28 to 2005 May 4, inferred the dynamical morphology of the active region. The 3-min narrowband signals detected over the sunspot and in the flare site are all well localised, which excludes their possible link with the instrumental artifacts, such as sidelobes of the image synthesis, and hence are natural. The 3-min narrowband maps of the active region, constructed with the use of PWF show the presence of extended V-shaped sources situated over the sunspot, with arms extended towards the flare site. We interpret these arms as evidence of the magnetic plasma channels that link the sunspot and the flare site by guiding magnetohydrodynamic waves. The 3-min periodicities of energy releases are then triggered by the 3-min oscillations leaking out from the sunspot along the magnetic structures.  On the basis of our findings, we deduce that the physical mechanism responsible for the relationship between 3-min sunspot oscillations and 3-min QPP in nearby flares can be as follows. The energy of 3-min oscillations leaks out of the sunspots in the form of field-aligned slow magnetoacoustic waves, which are often seen as compressible variations in the EUV radiation in the magnetic fan structures over sunspots (e.g., De Moortel 2006). In our study, these waves are seen as a 3-min modulation of the microwave radiation. The spatial distribution of the emission highlights the waveguiding plasma structures, which are the V-shaped microwave sources found in both flares discussed above. Because of the curvature of the magnetic field lines, the centrifugal force associated with the periodic longitudinal field-aligned wave motions produce periodic transverse kink-like perturbations of the magnetic structures (Zaitsev \\& Stepanov 1989). The induced transverse motions are fast magnetoacoustic waves that can carry energy and information across the magnetic field. These kink waves can trigger flaring energy releases (e.g., by the mechanism proposed by Nakariakov et al. 2006) provided that the waveguiding channel is situated close to the magnetic null-point, while is not necessarily linked magnetically. The modulation depth of the flaring light curves can be significantly stronger than in the modulating signal. Hence, the leaking 3-min wave can either trigger the energy release in a form of a short aperiodic spike or spike-like pulse, as seen on 2005 April 30, or lead to periodic triggering (or modulation) of energy releases in longer duration bursts, as seen on 2005 April 28 and 2005 May 4. In the first case, the triggered energy release uses up the energy stored in the magnetic configuration, and the next maximum in the wave cannot trigger another energy release. In the second case, either not all stored energy is liberated in the previous releases, or the next period of the triggering wave causes the energy release at another spatial location (see the discussion in Nakariakov et al. 2006).  In both analysed bursts, there is observational evidence that before the flares the energy of 3-min oscillations in the sunspot is enhanced significantly. In both cases, the amplitude of 3-min wave trains was highest just before the onset of the burst. This provides interesting perspectives on the use of the increase in the power of 3-min oscillations just before the flare as a flare precursor. However, the relationship between the amplitude of 3-min oscillations in the sunspot and energy releases nearby, requires statistical proof, and should be subject to a dedicated study.  A possible interpretation of the preferential appearance of the spike-like events when the analysed active region is in the vicinity of the central meridian, mentioned in Sect.~\\ref{sec:obs}, can be interpreted as either just the time coincidence with the appearance of certain physical conditions for the energy releases (e.g., the emergence of a new loop or arcade of loops from the sunspot to the burst source) or the preferential observability conditions (e.g., connected with the line-of-sight angle). The latter issue could be indicative of the similarity between the observed pulses and well-known microwave spikes: Altyntsev et al. (1996) found out that the spatial size of the microwave spike source is systematically larger in the vicinity of the limb. This was interpreted in terms of the scattering of the microwave emission across the coronal plasma. Hence, some events situated close the limb can be of lower intensity due to the scattering and hence be below the detection threshold and missed by observations. This may explain the preferential appearance of the short pulses in the vicinity of the central meridian."
    },
    "1005/1005.3077_arXiv.txt": {
        "abstract": "We conducted the observational tests of a phase correction scheme for the Atacama Compact Array (ACA) of the Atacama Large Millimeter and submillimeter Array (ALMA) using the Submillimeter Array (SMA). Interferometers at millimeter- and submillimeter-wave are highly affected by the refraction induced by water vapor in the troposphere, which results as phase fluctuations. The ACA is planning to compensate the atmospheric phase fluctuations using the phase information of the outermost antennas with interpolating to the inner antennas by creating a phase screen. The interpolation and extrapolation phase correction schemes using phase screens are tested with the SMA to study how effective these schemes are. We produce a plane of a wavefront (phase screen) from the phase information of three antennas for each integration, and this phase screen is used for the interpolation and extrapolation of the phases of inner and outer antennas, respectively. The interpolation scheme obtains apparently improved results, suggesting that the ACA phase correction scheme will work well. On the other hand, the extrapolation scheme often does not improve the results. After the extrapolation, unexpectedly large phase fluctuations show up to the antennas at the distance of $\\sim140$~m away from the center of the three reference antennas. These direction vectors are almost perpendicular to the wind direction, suggesting that the phase fluctuations can be well explained by the frozen phase screen. ",
        "introduction": "\\label{sect-intro}  The Atacama Large Millimeter and submillimeter Array (ALMA), the largest millimeter and submillimeter interferometer ever built, is currently under construction in the northern Chile with the collaboration between East Asia, Europe, and North America \\citep{woo09}. The ALMA is composed of up to eighty high-precision antennas, located on the Chajnantor plain of the Chilean Andes in the District of San Pedro de Atacama, 5000~m above sea level. The ALMA covers the wavelength range from 0.3~mm to 9~mm with an angular resolution of up to \\timeform{0''.004}. The Atacama Compact Array (ACA) is designed to improve the short baseline coverage of the ALMA, especially for the observations of extended structures at submillimeter wavelength \\citep{igu09}. The ACA consists of four 12~m telescopes to obtain the single-dish data and twelve 7~m telescopes to obtain short baseline interferometric data.  However, ground-base astronomical observations in the millimeter and submillimeter ranges are strongly affected by the fluctuation of the tropospheric water vapor distribution (e.g., \\cite{tho01}). Therefore, the correction of phase fluctuations due to the spatial and temporal variations of the tropospheric water vapor content is extremely important in millimeter and submillimeter interferometry. Several kinds of techniques have been proposed and performed for reducing phase fluctuations in millimeter and submillimeter interferometry (see \\cite{car99} for a review), such as fast switching phase calibration \\citep{hol92,hol95a,car97,mor00}, paired array phase calibration \\citep{hol92,car96,asa96,asa98}, and radiometric phase calibration \\citep{lay97,car98,car99,mar98,del00,wie01,mat02}. Although the ALMA site is one of the best sites for the millimeter and submillimeter astronomy \\citep{mat98,mat99,pai00,mat03}, atmosphere in this area is often affected by the phase fluctuations due to water vapor \\citep{rad96,hol97,but01}. The phase correction methods for the ALMA and the ACA are therefore highly needed to be considered and tested.  \\citet{asa05} conducted a series of simulations of a phase correction scheme for the ACA using water vapor radiometers (WVRs). The WVRs can measure the tropospheric water vapor content directly with observing a water vapor line in centimeter or millimeter wavelength. In the proposed ACA phase calibration scheme, the WVRs are attached to the 12~m antennas at the four corners of the twelve 7~m array (hereafter we call this as the reference rectangle). The changes of the tropospheric water vapor content aloft measured with the WVRs is transferred into the excess path lengths of the arriving radio waves. The excess path lengths measured at the four corners of the reference rectangle are then fitted to a simple two-dimensional slope or a screen. Then the phases of the antennas inside the reference rectangle can be compensated and calibrated. Their simulation succeeded to compensate the atmospheric phase well, which strongly support the use of the proposed phase correction scheme.  To confirm this simulation study observationally and discuss further, we performed the proposed phase correction scheme for the ACA using the Submillimeter Array (SMA; \\cite{ho04}). Here we present measurements with the SMA at 230~GHz, analyze the datasets under the proposed scheme, and discuss the results of the corrected phase variations. Our experiment is to clarify how effectively the proposed compensation scheme works under the conditions of the real atmosphere. We construct a reference triangle composed of three antennas, and make a flat phase plane or a screen with observing a strong point source (section~\\ref{sect-mdr}). The phases of antennas inside the reference triangle can be interpolated, while the phases of antennas outside can be extrapolated. We then compare the observed and the predicted phases of the point source. Standard deviations and temporal structure functions of the observed (uncorrected) and the corrected phases are also compared (section~\\ref{sect-res}). Finally, we discuss the usefulness of the interpolation and extrapolation phase correction schemes using a phase screen, and also discuss the validity of the ``frozen-flow'' hypothesis (section~\\ref{sect-dis}). ",
        "conclusions": "We performed an interferometric phase correction with the interpolation or extrapolation of the phase screen defined by three reference antennas using the SMA. This interpolation method is proposed for the ACA in the ALMA.  According to the comparisons of the standard deviations of corrected and uncorrected phases, relations between the phase standard deviation and the distance from the center of the reference triangle, and the temporal structure functions of the rms phase, the interpolation scheme improves phase fluctuation while the extrapolation scheme does not. This result can be explained by the boundary conditions of phase in these schemes; in case of the interpolation scheme, the phase corrected antenna is inside the triangle of three reference antennas, so the phase inside the triangle can be well defined (more known boundary conditions, more precisions or less phase errors and deviations). The extrapolation scheme, on the other hand, only has partial boundary conditions, and therefore less precision.  In the extrapolation scheme results, there is a sudden large phase fluctuation at the distance from the center of the reference triangle of about 140~m. According to the meteorological parameters on those observing dates and the antenna configurations, this ``140~m phase fluctuation'' is occurring only at the antennas located from the center of the reference triangle perpendicular to the wind direction. This ``140~m phase fluctuation'' can be explained by the frozen-flow model.  Although there are some differences between the configuration of our experiments and that in the proposed phase correction scheme for the ACA, our results based on the actual observations and the simulation results done by \\citet{asa05} promise the success of the phase correction for the ACA.   \\bigskip  We would like to thank Kazushi Sakamoto and Koh-Ichiro Morita for fruitful discussion. The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica."
    },
    "1005/1005.1072_arXiv.txt": {
        "abstract": "{We set out to determine the ratio, \\qir, of rest-frame 8--1000-$\\mu$m flux, \\sir, to monochromatic radio flux, \\s1.4, for galaxies selected at far-infrared (-IR) and radio wavelengths, to search for signs that the ratio evolves with redshift, luminosity or dust temperature, \\td, and to identify any far-IR-bright outliers -- useful laboratories for exploring why the far-IR/radio correlation (FIRRC) is generally so tight when the prevailing theory suggests variations are almost inevitable. We use flux-limited 250-$\\mu$m and 1.4-GHz samples, obtained using \\herschel\\ and the Very Large Array (VLA) in GOODS-North (-N).  We determine bolometric IR output using ten bands spanning $\\lambda_{\\rm obs}= 24-1250\\,\\mu$m, exploiting data from PACS and SPIRE (PEP; HerMES), as well as {\\it Spitzer}, SCUBA, AzTEC and MAMBO. We also explore the properties of an \\lir-matched sample, designed to reveal evolution of \\qir\\ with redshift, spanning log \\lir\\ = 11--12\\,L$_{\\odot}$ and $z=0-2$, by stacking into the radio and far-IR images. For 1.4-GHz-selected galaxies in GOODS-N, we see tentative evidence of a break in the flux ratio, \\qir, at \\l1.4\\ $\\sim 10^{22.7}$\\,W\\,Hz$^{-1}$, where active galactic nuclei (AGN) are starting to dominate the radio power density, and of weaker correlations with redshift and \\td. From our 250-$\\mu$m-selected sample we identify a small number of far-IR-bright outliers, and see trends of \\qir\\ with \\l1.4, \\lir, \\td\\ and redshift, noting that some of these are inter-related.  For our \\lir-matched sample, there is no evidence that \\qir\\ changes significantly as we move back into the epoch of galaxy formation: we find \\qir\\ $\\propto (1+z)^{\\gamma}$, where $\\gamma=-0.04\\pm 0.03$ at $z=0-2$; however, discounting the least reliable data at $z<0.5$ we find $\\gamma = -0.26\\pm 0.07$, modest evolution which may be related to the radio background seen by ARCADE\\,2, perhaps driven by $<$10-$\\mu$Jy radio activity amongst ordinary star-forming galaxies at $z>1$.} ",
        "introduction": "For samples of local galaxies -- on galactic and $\\sim$100-pc scales -- there is a good correlation between far-IR and radio emission \\citep{dejong85, helou85, condon91, yun01}. The correlation spans many orders of magnitude in luminosity, gas surface density and photon, cosmic-ray and magnetic energy density, and arises because the far-IR and radio wavelength regimes share a common link with luminous, massive stars and their end products -- dust, supernovae (SNe) and cosmic rays. In the simplest models \\citep[dubbed `calorimetry' -- e.g.][]{volk89, lvx96}, dust absorbs all of the ultraviolet radiation from massive stars, re-radiating this energy in the far-IR, and when those massive stars explode as SNe they generate cosmic-ray electrons which lose all their energy in the radio regime, mainly via synchrotron emission. A balance is thereby achieved between far-IR and radio emission, assuming that the starburst timescale is sufficiently long ($>$10$^7$\\,yr).  Traditionally, \\lir\\ and \\lrad\\ are both employed to determine star-formation rates, and the far-IR/radio flux density ratio has been useful when estimating the redshift or \\td\\ of a distant starburst, or when defining samples of AGN \\citep{condon92, cy99, ivison02, bell03, chapman05, donley05}, or probing magnetic field strength \\citep{thompson06}.  For these reasons, and because of recent observational advances at both far-IR and radio wavelengths, there has been a deluge of FIRRC-related work recently, exploring why the correlation exists and whether it continues to hold at progressively larger look-back times \\citep{garrett02, appleton04, ibar08, seymour09, ivison10, sargent10}. Prevailing theory \\citep[e.g.][]{lacki10} suggests that variations in the far-IR/radio flux ratio should be virtually unavoidable and that the FIRRC thus arises due to a mysterious combination of effects involving bremsstrahlung, inverse Compton cooling, ionisation and the relative fractions of primary/secondary cosmic-ray electrons/protons, as well as the critical synchrotron frequency.  Aside from the modelling work of \\citeauthor{lacki10}, recent advances in this field have included the use of luminosity-matched samples (between high and low redshift) to better probe evolution with look-back time \\citep{sargent10b} and the use of measurements spanning the far-IR and radio wavebands to avoid assumptions relating to $k$ corrections \\citep{ivison10}, although \\citet{calzetti10} have argued that bands beyond 24\\,$\\mu$m contain a contribution from dust heated by stars from previous episodes of star formation and so we might not necessarily expect the correlation to improve. In this paper we introduce flux-limited 250-$\\mu$m- and 1.4-GHz-selected samples of galaxies from \\herschel\\ and the VLA, as well as a luminosity-matched sample selected at 24\\,$\\mu$m, spanning $z=0-2$, and determine their spectral energy distributions (SEDs) spanning the entire far-IR spectral region.  We then investigate the FIRRC from the perspectives of the 24-, 250-$\\mu$m- and radio-selected samples. ",
        "conclusions": "\\begin{figure} \\centerline{\\psfig{figure=14552fg2.ps,width=3.3in,angle=270}} \\caption{\\qir\\ versus \\sir\\ for those 250-$\\mu$m-selected galaxies (sample~1) with secure, unambiguous radio ids.  Those without plausible radio ids are plotted as stars. The dashed line is the median, \\qir\\ = 2.40; the shaded region represents $\\pm$2$\\sigma_q$ ($\\sigma_q=0.24$).} \\label{q250} \\end{figure}  \\begin{figure} \\psfig{figure=14552fg3.ps,width=3.45in,angle=0} \\caption{\\qir\\ versus redshift for our radio-selected galaxies (sample~2), in five bins of $K$-corrected \\l1.4, plus the full sample. Values of \\qir\\ for sample 3 are shown (as circles) for comparison.} \\label{qlrad} \\end{figure}  \\begin{figure} \\psfig{figure=14552fg4.ps,width=3.55in,angle=0} \\caption{Median \\qir\\ versus \\l1.4. The local luminosity functions of starbursts and AGN are shown \\citep{ms07}.} \\label{ql1400} \\end{figure}  \\qir\\ as utilised here is the logarithmic ratio of the rest-frame 8--1000-$\\mu$m flux, \\sir, and the 1.4-GHz flux density, $S_{\\rm 1.4GHz}$, such that \\qir\\ = log$_{10}$ [(\\sir/$3.75\\times 10^{12}$\\,W\\,m$^{-2}$)/($S_{\\rm 1.4GHz}$/W\\,m$^{-2}$\\,Hz$^{-1}$)], where $S_{\\rm 1.4GHz}$ is $k$-corrected assuming $S_{\\nu}\\propto \\nu^{\\alpha}$, with $\\alpha=-0.8$.  We begin with sample (1),  those selected at 250\\,$\\mu$m: \\qir\\ is not a strong function of \\sir\\ (Fig.~\\ref{q250}), nor of \\s1.4. We see no evidence of contamination by radio-loud AGN, consistent with the findings of \\citet{yun01}. Some galaxies stand out as potentially far-IR-bright: these include the three galaxies without plausible radio ids, two of which are detected at 70 and/or 160\\,$\\mu$m, so are likely at low redshift with their radio emission resolved away.  Only 39/65 sources with unambiguous radio ids have redshifts (20 photometric, 19 spectroscopic; $\\left\\langle z \\right\\rangle = 0.98$; interquartile $z$ = 0.46--1.52, similar to sample~2).  Nevertheless, this sub-sample allows us to explore correlations between \\qir\\ and luminosity, redshift and \\td. We find significant ($>$95\\% confidence -- Table~\\ref{trends}) trends for lower \\qir\\ amongst the most radio- and far-IR-luminous galaxies, and the warmest and most distant, though these parameters are likely inter-related. The dependence of \\qir\\ on \\l1.4\\ is the strongest and likely reflects the influence of low-radio-power AGN, of which more later; that of \\qir\\ on \\lir\\ is more puzzling, perhaps reflecting the dependence of \\lir\\ on redshift and/or \\td\\ \\citep[e.g.][]{chapman05}, or selection effects (since this trend is not seen for sample 2 -- see Table~\\ref{trends}).  \\begin{table} \\caption{Trends.} \\label{trends} \\centering \\begin{tabular}{l c c} \\hline\\hline \\qir\\ trend& Spear-    &Signifi-\\\\ & man $\\rho$&-cance\\\\ \\hline \\multicolumn{3}{l}{Sample 1 (250-$\\mu$m-selected galaxies with redshifts):}\\\\ (5.02$\\pm$0.18) -- (0.105$\\pm$0.008)\\,log\\,\\l1.4 &$-0.48$&99.8\\%\\\\ (6.09$\\pm$0.33) -- (0.092$\\pm$0.008)\\,log\\,\\lir  &$-0.32$&95.6\\%\\\\ (2.61$\\pm$0.02) -- (0.081$\\pm$0.007)\\,(1+$z$)    &$-0.33$&96.0\\%\\\\ (2.76$\\pm$0.03) -- (0.008$\\pm$0.001)\\,\\td\\       &$-0.33$&96.1\\%\\\\ \\hline \\multicolumn{3}{l}{Sample 2 (radio-selected galaxies with redshifts):}\\\\ (4.92$\\pm$0.21) -- (0.101$\\pm$0.009)\\,log\\,\\l1.4 &$-0.27$&99.9\\%\\\\ (2.74$\\pm$0.35) -- (0.007$\\pm$0.009)\\,log\\,\\lir  &$+0.07$&69.1\\%\\\\ (2.55$\\pm$0.02) -- (0.047$\\pm$0.010)\\,(1+$z$)    &$-0.15$&96.6\\%\\\\ (2.60$\\pm$0.02) -- (0.002$\\pm$0.001)\\,\\td\\       &$-0.16$&89.2\\%\\\\ \\hline \\end{tabular} \\end{table}  Fig.~\\ref{qlrad} shows \\qir\\ versus redshift for our radio-selected galaxies (sample 2), split into five log-spaced bins of \\l1.4. Does \\qir\\ evolve with redshift? One might conclude that it does, based on the bottom panel of Fig.~\\ref{qlrad}, where \\qir\\ $\\propto (1+z)^{\\gamma}$, with $\\gamma=-0.05\\pm 0.01$ (Table~\\ref{trends}). However, we must be aware of some strong selection effects which make this evidence unreliable: radio emission can be due to an AGN and several radio-loud objects with low values of \\qir\\ are obvious in Fig.~\\ref{qlrad}. Such AGN are more common at $z\\sim 2$ than today \\citep[e.g.][]{wall05}; moreover, radio emission from faint starbursts (with $\\alpha=-0.8$, although see \\citealt{ibar10}) becomes more difficult to detect at higher redshifts, such that the fraction of radio-loud AGN in a flux-limited sample will rise, driving down \\qir.  Indeed, Fig.~\\ref{ql1400} shows tentative evidence of a break in $\\left\\langle q_{\\rm IR} \\right\\rangle$ at \\l1.4\\ $\\sim10^{22.7}$\\,W\\,Hz$^{-1}$. One might also expect radio-loud objects (those with low \\qir) to contain warmer, AGN-heated dust, giving rise to the weak trend (89.2\\% confidence -- Table~\\ref{trends}) of decreasing \\qir\\ with increasing \\td.  \\begin{figure} \\centerline{\\psfig{figure=14552fg5.ps,width=3.4in,angle=270}} \\caption{\\lir\\ (dots; left axis) and \\qir\\ (red circles; right axis) -- the former determined via the models of \\citet{ce01} -- versus redshift for our \\lir-matched sample.  The luminosity bounds and redshift bins (dashed lines), the number of galaxies in each bin and their predicted (measured) $\\left\\langle {\\rm log} L_{\\rm IR}\\right\\rangle$ and measured $\\left\\langle q_{\\rm IR}\\right\\rangle$ are all shown.  The shaded area represents a $\\pm 1\\sigma$ prediction for \\qir\\ \\citep{swinbank08, ivison10}.} \\label{lirz} \\end{figure}  Finally, we turn to our \\lir-matched galaxies (sample 3), illustrated in Fig.~\\ref{lirz}.  The $\\delta z=0.5$ bins provide significant numbers of objects at near-constant \\lir\\ spanning $z=0-2$. As well as being matched in \\lir, there is another key difference between our new sample and that used by \\citet{ivison10}: although the new sample is based initially on a flux-limited 24-$\\mu$m catalogue, the final selection is based on \\lir, with model-dependent extrapolations from the mid-IR (accurate to $\\ls2\\times$ across all bins -- Fig.~\\ref{lirz}). This should lead to less contamination by AGN at the blue end of the rest-frame 8--1000-$\\mu$m band, where the relative contribution to \\sir\\ can be substantial \\citep[Figure~11 --][]{ivison10}. Using our new sample, there is no strong evidence that \\qir\\ changes as we move back into the epoch of galaxy formation at $z\\sim 2$, with $\\gamma=-0.04\\pm 0.03$ where \\qir\\ $\\propto (1+z)^{\\gamma}$, consistent with the findings of \\citet{sargent10b}. If we discount the $z<0.5$ data, which comprise only 16 galaxies which are not well matched in \\lir\\ to the higher redshift bins, we find $\\gamma = -0.26\\pm 0.07$. This is similar to the $\\gamma = -0.15\\pm 0.03$ found by \\citet{ivison10} who noted reports that evolution in \\qir\\ could be related to the radio background seen by ARCADE\\,2 \\citep{fixsen10, seiffert10}. Our sample, with $\\left\\langle S_{\\rm 1.4GHz} \\right\\rangle \\ls 10\\,\\mu$Jy at $z\\gs 1$, is consistent with the idea that evolution of the FIRRC might be driven by $<$10-$\\mu$Jy radio activity amongst ordinary star-forming galaxies at $z>1$ \\citep{singal10}."
    },
    "1005/1005.1591_arXiv.txt": {
        "abstract": "{It has been shown that by means of different physical mechanisms the expansion of H\\,{\\sc{ii}} regions can trigger the formation of new stars of all masses. This process may be important to the formation of massive stars but has never been quantified in the Galaxy. } {We use $\\it{Herschel}$-PACS and -SPIRE images from the $\\it{Herschel}$ Infrared survey of the Galactic plane, Hi-GAL, to perform this study.} {We combine the $\\it{Spitzer}$-GLIMPSE and -MIPSGAL, radio-continuum and sub-millimeter surveys such as ATLASGAL with Hi-GAL to study Young Stellar Objects (YSOs) observed towards Galactic \\HII\\ regions. We select a representative \\HII\\ region, N49, located in the field centered on $l$=30$\\degr$ observed as part of the Hi-GAL Science Demonstration Phase, to demonstrate the importance Hi-GAL will have to this field of research. } {Hi-GAL PACS and SPIRE images reveal a new population of embedded young stars, coincident with bright ATLASGAL condensations. The Hi-GAL images also allow us, for the first time, to constrain the physical properties of the newly formed stars by means of fits to their spectral energy distribution. Massive young stellar objects are observed at the borders of the N49 region and represent second generation massive stars whose formation has been triggered by the expansion of the ionized region.} {The first Hi-GAL images obtained using PACS and SPIRE have demonstrated the capability to investigate star formation triggered by \\HII\\ regions. With radio, submillimeter, and shorter wavelength infrared data from other surveys, the Hi-GAL images reveal young massive star-forming clumps surrounding the perimeter of the N49 \\HII\\ generated bubble. Hi-GAL enables us to detect a population of young stars at different evolutionary stages, cold condensations only being detected in the SPIRE wavelength range. The far IR coverage of Hi-GAL strongly constrains the physical properties of the YSOs. The large and unbiased spatial coverage of this survey offers us a unique opportunity to lead, for the first time, a global study of star formation triggered by \\HII\\ regions in our Galaxy. } ",
        "introduction": "Ionized (H\\,{\\sc{ii}}) regions are known to trigger the formation of stars by means of various physical mechanisms (see Elmegreen~\\cite{elm98} and Deharveng et al.~\\cite{deh05} for a review). Several H\\,{\\sc{ii}} regions have been studied individually in the context of triggered star formation, focusing on the associated neutral material and the young stellar population (Zavagno et al.~\\cite{zav06}; Deharveng et al.~\\cite{deh09}; Pomar\\`es et al.~\\cite{pom09}; Bieging et al.~\\cite{bie09}). These studies have shown that the expansion of H\\,{\\sc{ii}} regions can trigger the formation of new stars of all masses. The $\\it{Spitzer}$-GLIMPSE survey of the Galactic plane (Benjamin et al.~\\cite{ben03}) detected nearly 600 bubbles (Churchwell et al. \\cite{chu06}). Deharveng et al. (\\cite{deh10}) selected a series of 102 ionized bubbles and studied the star formation in their surroundings using $\\it{Spitzer}$-GLIMPSE and MIPSGAL (Carey et al. \\cite{car09}), radio (MAGPIS; Helfand et al. \\cite{hel06} and VGPS; Stil et al. \\cite{sti06}), and ATLASGAL (Schuller et al. \\cite{sch09}) data. They show that 86\\% of these bubbles enclose H\\,{\\sc{ii}} regions, and that more than 20\\% of 64 bubbles (for which the ATLASGAL angular resolution is sufficient to resolve the spatial distribution of cold dust) show massive star formation on their borders. This indicates that triggering is important in the creation of massive stars and that hot photodissociation regions (PDRs) are a good place to study the earliest phases of massive-star formation.  Our long-term aim is to use Hi-GAL (Molinari et al.~\\cite{mol10}), combined with existing infrared, submillimeter, and radio surveys, to study the influence of H\\,{\\sc{ii}} regions on triggering the formation of stars in our Galaxy. Hi-GAL's extended wavelength coverage towards the far-infrared and its unprecedented sensitivity offer a unique opportunity to detect an embedded population of young sources that are not detected at shorter wavelengths. This allows us to observe intermediate and high-mass YSOs over the complete range of evolutionary stages. The unprecedented resolution of Hi-GAL also offers the opportunity to accurately characterize the physical nature of the sources by means of a detailed fit to their spectral energy distribution (SED).  In this Letter we study the bubble-shaped ionized region, N49, from the Churchwell et al.~(\\cite{chu06}) catalogue to illustrate the purpose of our project. The N49 bubble was studied by Watson et al.~(\\cite{wat08}) and by Deharveng et al.~(\\cite{deh10}). However, these studies had no information in the 70--500\\,$\\mu$m range. This information, obtained with the PACS and SPIRE data presented here, allows us to discuss the star formation of this region in detail. ",
        "conclusions": "We have presented the first PACS and SPIRE images from Hi-GAL of the bubble-shaped Galactic \\HII\\ region N49. This region is used as an illustration of the study dedicated to the star formation triggered by Galactic \\HII\\ regions that we plan to lead for the whole survey, combining the Hi-GAL results with other infrared and radio surveys of the Galactic plane. We have shown that: \\begin{itemize} \\item{The Hi-GAL SPIRE and PACS images allow us to study the distribution of young sources towards N49. The far-IR fluxes have been measured and strongly constrain the spectral energy distribution of these sources. This allows us to characterize their properties. } \\item{The PACS images reveal the existence of red young stellar objects towards two ATLASGAL condensations, sources that had not been previously detected at shorter wavelengths. The ultracompact \\HII\\ region is not clearly seen in the $\\it{Herschel}$ range. The bright YSO\\#3 Watson et al. observed to be coincident with condensation 1 dominates the emission at longer wavelengths.} \\item{SED fits for the 3 sources detected by $\\it{Herschel}$ towards millimeter condensations using the Robitaille et al.~(\\cite{rob07}) model show that these sources are young and massive. However, their age has not been constrained and other indicators are needed to refine the discussion of star formation history in this region.} \\item{Five massive stars are forming in the N49 PDR. The high star formation efficiency in N49 may be due to the presence of winds from the first generation massive star.} \\end{itemize} The study of N49 using a multiwavelength approach shows that Hi-GAL enables measurements for the crucial far-IR range to be made that are essential to infer the properties of the YSOs. Hi-GAL clearly provides important insight into star formation triggered by expanding \\HII\\ regions. Seventy-six \\HII\\ regions of all shapes are detected in the $l$=30$\\degr$ field and a higher density of Hi-GAL sources is clearly observed towards these regions, indicating that star formation triggered by \\HII\\ regions may be an important process."
    },
    "1005/1005.1244_arXiv.txt": {
        "abstract": "{It is difficult to determine masses and test formation models for brown dwarfs, because they are always above the main sequence, so that there is a degeneracy between mass and age. However, for brown dwarf companions to normal stars, such determinations may be possible, because one can know the distance and age of the primary star. As a result, brown dwarf companions are well-suited to testing formation theories and atmosphere models. } {With more adaptive optics images available, we aim at detecting orbital motion for the first time in the system TWA 5 A+B. } {We measured separation and position angle between TWA 5 A and B in each high-resolution image available and followed their change in time, because B should orbit around A. The astrometric measurement precision is about one milli arc sec. } {With ten year difference in epoch, we can clearly detect orbital motion of B around A, a decrease in separation by $\\sim 0.0054^{\\prime \\prime}$ per year and a decrease in position angle by $\\sim 0.26^{\\circ}$ per year. } {TWA 5 B is a brown dwarf with $\\sim 25$ Jupiter masses (Neuh\\\"auser et al. 2000), but having large error bars (4 to 145 Jupiter masses, Neuh\\\"auser et al. 2009). Given its large projected separation from the primary star, $\\sim 86$ AU, and its young age ($\\sim 10$ Myrs), it has probably formed star-like, and would then be a brown dwarf companion. Given the relatively large changes in separation and position angle between TWA 5 A and B, we can conclude that they orbit around each other on an eccentric orbit. Some evidence is found for a curvature in the orbital motion of B around A - most consistent with an elliptic (e=0.45) orbit. Residuals around the best-fit ellipse are detected and show a small-amplitude ($\\sim 18$ mas) periodic sinusoid with $\\sim 5.7$ yr period, i.e., fully consistent with the orbit of the inner close pair TWA 5 Aa+b. Measuring these residuals caused by the photocenter wobble - even in unresolved images - can yield the total mass of the inner pair, so can test theoretical pre-main sequence models. }  \\titlerunning{Orbital motion of TWA 5 B} ",
        "introduction": "The star TWA 5 is one of the five original members of the TW Hya association (TWA), a group of 5 to 12 Myr young stars (Kastner et al. 1997), where no gas clouds are left from the star formation process (Tachihara et al. 2009); see Torres et al. (2008) for a recent review on TWA. TWA 5 is an M1.5 weak-line T Tauri star (Webb et al. 1999) with variable H$\\alpha$ emission, hence still ongoing accretion (Mohanty et al. 2003). The central star itself is either a very close ($\\le 66$ milli arc sec or mas) binary (Konopacky et al. 2007, henceforth K07) or even triple (Torres et al. 2003). The close inner pair TWA 5 Aa+b has a total dynamical mass of $0.71 \\pm 0.14$~M$_{\\odot}$ (assuming 44 pc as distance) and an orbital period of $5.94 \\pm 0.09$ years (K07). The wide companion TWA 5 B was originally discovered by Webb et al. (1999) and Lowrance et al. (1999) and confirmed as co-moving with TWA 5 A by Neuh\\\"auser et al. (2000).  The spectral type of TWA 5 B is M8-9 (Webb et al. 1999, Lowrance et al. 1999, Neuh\\\"auser et al. 2000, Mohanty et al. 2003). The mass of the companion is between 15 and 40 Jupiter masses just from temperature, luminosity, and theoretical hot-start model tracks (Neuh\\\"auser et al. 2000). The mass lies anywhere between 4 and 145 Jupiter masses, if calculated from temperature ($2800 \\pm 100$ K), luminosity ($\\log(L_{bol}/L_{\\odot}) = -2.62 \\pm 0.30$ at $44 \\pm 4$ pc), and gravity ($\\log g = 4.0 \\pm 0.5$), as obtained by comparison of a Sinfoni K-band spectrum with Drift-Phoenix model atmospheres (Neuh\\\"auser et al. 2009).  The system TWA 5 A+B was observed by several teams with ground-based adaptive optics (AO) and/or the Hubble Space Telescope (HST). We obtained two more recent images, so that we can now investigate possible orbital motion of B around A with a 10 year difference in epoch (first preliminary results in Schmidt et al. 2008). We can then also try to detect the orbital motion of Ab and Aa around each other as residuals of the much longer orbit of B around A due to a periodic wobble of the photocenter of the close Aa+Ab pair. ",
        "conclusions": ""
    },
    "1005/1005.5301_arXiv.txt": {
        "abstract": "A high-order scheme for direct numerical simulations of turbulent combustion is discussed. Its implementation in the massively parallel and publicly available {\\sc Pencil Code} is validated with the focus on hydrogen combustion. Ignition delay times (0D) and laminar flame velocities (1D) are calculated and compared with results from the commercially available Chemkin code. The scheme is verified to be fifth order in space. Upon doubling the resolution, a 32-fold increase in the accuracy of the flame front is demonstrated. Finally, also turbulent and spherical flame front velocities are calculated and the implementation of the non-reflecting so-called Navier-Stokes Characteristic Boundary Condition is validated in all three directions. ",
        "introduction": "Modeling of turbulence is one of the largest research areas within flow mechanics. Turbulent combustion inherits all the properties of non-reacting turbulent flow. The most important addition is linked to the highly nonlinear reaction processes, and models for this are called combustion models. Two additional challenges in turbulent combustion are the very sharp changes in density and differential diffusion of mass and heat.  For combustion processes it is crucial to be able to simulate the mixing of the combustible species correctly. Traditionally this has been done by means of mixing models in Reynolds Averaged Navier Stokes (RANS) codes by combining, e.g., the $k$-$\\epsilon$ turbulence model and the eddy dissipation concept (or EDC) ``mixing'' model \\citep{MH76,M81,M89}, or in Large Eddy Simulation (LES) \\citep{PV05} where a sub-grid model is used both for the turbulence and for the scalar mixing. There are however major and still unresolved problems related to modelling of what happens on the very smallest scales with these methods.  Several RANS codes with detailed chemistry are commercially available \\citep{c10,f10}, and there are a huge number of these codes found as in-house codes at different academic institutions and in many industrial departments around the world. There are also freely available open-source RANS codes with detailed chemistry \\citep{o10}. The reason for the popularity of RANS is its low demand on computational resources. Because of this RANS has, for decades, been the most used type of code for industrial purposes.  Nevertheless, also LES has increased in popularity during the last years, and this has led, for example, to the inclusion of a LES module in \\cite{f10}. Most LES codes for combustion today are, however, in-house codes owned by different academic institutions.  The most accurate way of simulating turbulent combustion is to use Direct Numerical Simulation (DNS)~\\citep{PV05} instead of RANS or LES. In DNS one resolves the full range of time and length scales of both the turbulence and the chemistry (using accurate high-order numerical methods for computational efficiency). The problem with DNS is however that it is very resource demanding, both on CPU-hours and memory.  In this paper we present the implementation of a detailed chemistry module in a finite-difference code \\citep{PC} for compressible hydrodynamic flows. The code advances the equations in non-conservative form. The degree of conservation of mass, momentum and energy can then be used to assess the accuracy of the solution. The code uses six-order centered finite differences. For turbulence calculation we normally use the RK3-2N scheme of \\cite{Wil80} for the time advancement \\citep{BD02}. This scheme is of Runge-Kutta type, third order, and it uses only two chunks of memory for each dependent variable. For hydrodynamic calculations, the lengths of the time step is calculated based on a number of constraints involving maximum values of velocity, viscosity, and other quantities on the right-hand sides of the evolution equations. In some cases we use instead a fifth-oder Runge-Kutta-Fehlberg scheme with an automatic adaptive time step, subject to the aforementioned hydrodynamic constraints. However, in many cases we found it advantageous to use a fixed time step whose length is estimated based on earlier trial runs with an automatically calculated time step.  On a typical processor, the cache memory between the CPU and the RAM is not big enough to hold full three-dimensional data arrays. Therefore, the {\\sc Pencil Code} has been designed to evaluate first all the terms on the right-hand sides of the evolution equations along a one-dimensional subset (pencil) before going to the next pencil. This implies that all derived quantities exist only along pencils. Only in exceptional cases do we allocate full three-dimensional arrays to keep derived quantities in memory. However, most of the time, multiple operations including the calculation of derivatives is performed without using intermediate storage.  As far as we are aware, no open source high-order DNS code with detailed chemistry is currently available. The amount of man-hours for implementing a fully parallelized DNS code with detailed chemistry is enormous. It is therefore now timely to make such a code available in the public domain and to encourage further development by a wider range of scientists. Here we describe the implementation of such a scheme in the {\\sc Pencil Code}, which is currently maintained under the Google Code subversion repository, \\url{http://pencil-code.googlecode.com/}. The code is highly modular and comes with a large selection of physics modules. It is portable to all commonly used architectures using Unix or Linux operating systems. The code is well documented and independent of external libraries and any third party licenses. All parts of the code, including the current chemistry implementation, is therefore explicitly open source code. In particular, there are no pre-compiled binary files. Consequently there are no licenses required for running any part of the code. It is therefore straightforward to download the full source code from the original subversion repository on google-code. The Message Passing Interface libraries are needed when running on multiple processors, but all parts of the code can also run on a single processor without these libraries. The integrity of the code is monitored through the automatic execution of a selection of test cases on various platforms at different sites. The detailed history of the code with about 14,000 revisions is accessible.  It should be emphasized that the use of high-order discretization is critical for optimizing the accuracy at a given resolution. Doubling the resolution of a 3D explicit code require 16 times more CPU time, but this increases the accuracy by a factor of 32. In fact, switching to a derivative module with a tenth order scheme is straightforward and not significantly more expensive. ",
        "conclusions": "In this paper we have presented a high-order public domain code for direct numerical simulation of compressible flows with detailed chemical reactions. The {\\sc Pencil Code} provides sixth-order spatial accuracy in the simple one-step reaction case, and fifth order accuracy in the case where upwinding for density advection is necessary. For validation purposes we compare our results with the Chemkin tool for 0D and 1D test problems, and show that they are in good agreement. Finally, we calculate the flame speed in 3D both in laminar and turbulent cases.  The code is well suited for considering also more complicated reaction schemes such as methane combustion. Furthermore, it is straightforward to consider the interaction with additional chemicals such as nitrogen and to follow the production of NOx gases. In particular, it is important to consider combustion in the presence of steam. This is well known to lead to a reduction of NOx gases. Combustion in the presence of more complicated boundary conditions involving, for example, smaller inlet geometries has also been considered. Some of these cases, including those with a turbulent inlet, are available among the many sample cases that come with the code. For the benefit of the community, it is advantageous if prospective contributers to the code ask one of the code owners listed on \\url{http://pencil-code.googlecode.com/} to obtain permission as a committer."
    },
    "1005/1005.2521_arXiv.txt": {
        "abstract": "{} { We identify a prominent absorption feature at 1115 GHz, detected in first HIFI spectra towards high-mass star-forming regions, and interpret its astrophysical origin. } { The characteristic hyperfine pattern of the \\hoplus{} ground-state rotational transition, and the lack of other known low-energy transitions in this frequency range, identifies the feature as \\hoplus{} absorption against the dust continuum background and allows us to derive the velocity profile of the absorbing gas. By comparing this velocity profile with velocity profiles of other tracers in the DR21 star-forming region, we constrain the frequency of the transition and the conditions for its formation. } { In DR21, the velocity distribution of \\hoplus{} matches that of the \\CII{} line at 158\\ ${\\mu}$m and of OH cm-wave absorption, both stemming from the hot and dense clump surfaces facing the \\HII{}-region and dynamically affected by the blister outflow. Diffuse foreground gas dominates the absorption towards Sgr~B2. The integrated intensity of the absorption line allows us to derive lower limits to the \\hoplus{} column density of $7.2\\times 10^{12}$~cm$^{-2}$ in NGC~6334, $2.3\\times 10^{13}$~cm$^{-2}$ in DR21, and $1.1\\times 10^{15}$~cm$^{-2}$ in Sgr~B2. } {} ",
        "introduction": "Oxidaniumyl or oxoniumyl \\citep[][]{IUPAC}, the reactive water cation, \\hoplus{}, plays a crucial role {\\fourthchanged in the chemical network describing the formation of oxygen-bearing molecules} in UV irradiated parts of molecular clouds \\citep{vDh_Black,Gerin}. It was identified at optical wavelengths in the tails of comets in the 1970's \\citep{Fehrenbach,H2O+_Kohoutek_1974a,H2O+_Kohoutek_1974b}, but its detection in the general interstellar medium has proven elusive.  We report a detection of the ground-state rotational transition of \\hoplus{} in some of the first spectra taken with the HIFI instrument \\citep{HIFI} on board the {\\it Herschel} Space Observatory \\citep{Herschel} during the performance verification campaign and early science observations. Section 2 briefly introduces the properties of the sources where \\hoplus{} was detected. Section 3 summarises the spectroscopic data of the molecule. The observations and the line identification are described in Sect. 4 {\\fourthchanged and in Sect. 5 we discuss} the physical properties of the \\hoplus{} absorption layer. ",
        "conclusions": "That \\hoplus{} shows up in absorption against the dust continuum implies that the excitation of the molecule must be colder than the dust. As a reactive ion (see the discussion  by \\citealt[][]{Black,Staeuber} for CO$^+$), \\hoplus{} is not expected to be in thermal equilibrium at the kinetic temperature of the gas. Its excitation reflects either the chemical formation process or the radiative coupling with the environment. From a single absorption line, {\\fourthchanged one can only} provide a lower limit to the \\hoplus{} column density, assuming a low excitation temperature where basically all \\hoplus{} resides in the ground state, which is applicable to temperatures well below the upper level energy of 53~K.  Table~\\ref{line_parameter} provides the integral over the optical depth of the hyperfine components in the low temperature limit. For the overall $J=3/2-1/2$ fine structure transition, we obtain a line integrated optical depth of $\\int \\tau d\\nu/N\\sub{H_2O^+}=4.70\\,10^{-13}$ km s$^{-1}$ cm$^2$ per molecule, resulting in a lower limit to the \\hoplus{} ground-state column densities of  $7.2\\times 10^{12}$~cm$^{-2}$ for NGC~6334, $2.3\\times 10^{13}$~cm$^{-2}$ for DR21, and $1.1\\times 10^{15}$ cm$^{-2}$ for Sgr~B2.  These values are lower limits not only because of to the low-temperature approximation, but also because they assume that the absorption occurs in front of the continuum source and {\\fourthchanged not within the dusty cloud}, where the line absorption is partially compensated by dust emission. There may also be additional amounts of \\hoplus{} in the para species that would not contribute to the 1115~GHz line. Altogether, the total \\hoplus{} column density could be much higher than the lower limits given here.    The excellent correlation between the \\hoplus{} profile and the OH absorption profile in DR21 indicates that both species occur in the same thin layer of hot gas \\citep{Jones} that directly faces the \\HII{} region at the blue-shifted blister outflow. There is no obvious correlation with the distributions of CO, H$_2$O, or HCO$^+$. For Sgr~B2, we can clearly identify absorption in multiple translucent foreground clouds. Their densities must be high enough to produce some molecular hydrogen, but low enough not to quickly destroy the \\hoplus{}. For NGC6334, the gas component producing the \\hoplus{} absorption remains unidentified.  {\\changed With the identification of H$_2$O$^+$ in the interstellar medium, we provide a {\\thirdchanged first step to quantifying} an important intermediate node in the oxygen chemical network, connecting OH$^+$ in diffuse clouds and at cloud boundaries, through H$_3$O$^+$, with water in denser and cooler cloud parts.} To obtain {\\fourthchanged an estimate for the total \\hoplus{} abundance, we need to measure the excitation temperature of \\hoplus{}.} Observations of additional transitions of \\hoplus{}, such as those at 742~GHz, are therefore essential."
    },
    "1005/1005.0148_arXiv.txt": {
        "abstract": "We constrain a stochastic background of primordial magnetic field (PMF) by its contribution to the angular power spectrum of cosmic microwave background anisotropies. We parametrize such stochastic background by a power-law spectrum with index $n_B$ and by its Gaussian smoothed amplitude $B_\\lambda$ on a comoving length $\\lambda$. We give an approximation for the spectra of the relevant correlators of the energy-momentum of the stochastic background of PMF for any $n_B$. By using the WMAP 7 year data in combination with ACBAR, BICEP and QUAD we obtain the constraint $B_{1 {\\rm Mpc}} < 5.0$ nG at $95\\%$ confidence level for a stochastic background of non-helical PMF. We discuss the relative importance of the scalar and vector contribution to CMB anisotropies in obtaining these constraints. We then forecast {\\sc Planck} capabilities in constraining $B_{1 {\\rm Mpc}}$. ",
        "introduction": "The origin of the large scale magnetic fields observed in galaxies and clusters of galaxies is an open issue of great importance in modern astrophysics (see \\cite{Widrow:2002ud} for a review). Primordial magnetic fields (PMF) generated in the early Universe could have been the seeds for large scale magnetic fields and have left an imprint in the anisotropy pattern of the cosmic microwave background (CMB). A primordial hypothesis for generating the seeds amplified afterwards by adiabatic compression and dynamo - cannot be discarded \\cite{Widrow:2002ud}, also in light of recent observations of strong magnetic fields in galaxies at high redshift \\cite{Bernet:2008qp,Wolfe:2008nk}.  PMF with a comoving amplitude of several nG can leave interesting imprints on CMB anisotropies. A stochastic background of PMF is modelled as a fully inhomogeneous component and its energy momentum tensor (EMT) - quadratic in the magnetic fields - is considered at the same footing as linear inhomogeneities in the other components and linear metric fluctuations. A stochastic background of PMF generates independent modes for all kinds of linear perturbations: there has been several studies for scalar \\cite{KL,yamazaki,KR,GK,FPP,PFP,SL,BoCa}, vector \\cite{SB,MKK,lewis} and tensor \\cite{DFK,MKK,CDK} perturbations. See Refs. \\cite{PFP,SL} for studies which take into account all the types of contributions. A stochastic background of PMF affects also the statistics of CMB anisotropies, and not only its power spectra: being quadratic in the magnetic field amplitude, the EMT is non-Gaussian distributed \\cite{BC} and therefore the bispectrum of CMB anisotropies can also be a useful probe \\cite{CFPR,SS}.  In our previous works \\cite{FPP,PFP} we have refined the computation of CMB anisotropies in presence of a stochastic background of PMF: we have computed the initial conditions for cosmological perturbations in the radiation era keeping into account only relativistic degrees of freedom and the correlators for the Fourier transforms of the EMT in presence of a sharp cut-off which mimics the damping scale due to viscosity \\cite{FPP,PFP} for few values of the spectral index $n_B$ of the stochastic background.  In this paper we use and extend the above theoretical description to derive the CMB constraints on a stochastic background of PMF which can be obtained by current and future data. We therefore use a modified version of \\texttt{CosmoMC} \\cite{cosmomc} connected with a modified version of \\texttt{CAMB} \\cite{camb} containing all the above features to constrain $B_\\lambda$ - the amplitude of the magnetic field smoothed over a comoving scale $\\lambda$ - and $n_B$ with the most recent compilation of CMB anisotropies data, therefore updating previous investigations \\cite{yamazaki,yamazakilast,shawlewis2}. We also forecast the {\\sc Planck} \\cite{bluebook} capabilities in constraining such a background of PMF.  Our paper is organized as follows. In Section II we present the theoretical set-up needed for deriving the CMB constraints. As new results we give an approximation for the PMF energy-momentum valid for any $n_B$, we discuss the correlator between the energy density and the Lorentz force and we extend the regular initial conditions for cosmological perturbations in the relativistic regime in presence of a stochastic background of PMF \\cite{PFP} including the correction due to matter, collecting the details in Appendix A, B and C, respectively. In Section III we discuss the constraints from WMAP 7 years data \\cite{jarosik,larson}, ACBAR \\cite{ACBAR}, BICEP \\cite{BICEP} and QUAD \\cite{QUAD} on a flat $\\Lambda$CDM model plus a stochastic background of PMF. We present the Planck capabilities in constraining $B_{\\lambda}$ and $n_B$ in Section IV and we summarize our results in Section V. ",
        "conclusions": "We studied the constraints on a stochastic background of PMF by current and forthcoming CMB data. In doing this, we have improved the theoretical CMB predictions in presence of a stochastic background of PMF. We gave approximations for the relevant components of the EMT for any spectral index $n_B$. We considered the correlation between $\\rho_B$ and $L_B$ and showed how the previous choice of total anti-correlation \\cite{FPP,PFP} was a rather good approximation, in particular for red values of $n_B$. We computed the initial conditions for magnetic scalar mode in presence of matter corrections showing how these corrections do not affect the results contrary to what claimed in \\cite{Kojima}.  On the basis of previous works, we have considered only the scalar and vector contribution, and by using their regular initial condition, we constrain $B_{1 {\\rm Mpc}} < 5.0$ nG and $n_B < - 0.12$ at $95\\%$ confidence level. with the most updated combination of CMB anisotropies. {\\sc Planck} will be able to constrain the spectrum of a stochastic background of PMF even further at the level of $2.7$ nG.  \\vspace{1cm}  \\begin{center} {\\bf Acknowledgements} \\end{center} We wish to thank C. Caprini, R. Durrer, J. Hamann and L. Hollenstein for useful comments. This work has been done in the framework of the {\\sc Planck} LFI activities and is partially supported by ASI contract {\\sc Planck} LFI activity of Phase E2. We acknowledge the use of the Legacy Archive for Microwave Background Data Analysis (LAMBDA). Support for LAMBDA is provided by the NASA Office of Space Science. \\begin{widetext}        \\begin{figure} \\includegraphics[scale=1]{n23scal.eps}\\includegraphics[scale=1]{nm27scal.eps}\\\\ \\includegraphics[scale=1]{n23L.eps}\\includegraphics[scale=1]{nm27L.eps}\\\\ \\includegraphics[scale=1]{n23V.eps}\\includegraphics[scale=1]{nm27V.eps} \\caption{Comparison of spectral fit (dashed line) and exact spectrum (solid line) for the rescaled magnetic energy density $|\\rho_B (k)|^2/U$ (top panel), scalar Lorentz force $|L_B (k)|^2/U$ (middle panel) and the vector part of the anisotropic stress $|\\Pi^{(V)} (k)|^2/(2 U)$ (bottom panel) as a function of $k/k_D$ with $U=(A^2 k_D^{2n_B+3})/(512 \\pi^4) $. In the left (right) column $n_B=2.3$ ($n_B=-2.7$) is displayed.} \\label{Comparison_fit_scal} \\end{figure}   \\appendix"
    },
    "1005/1005.5071_arXiv.txt": {
        "abstract": "Nucleosynthetic yield predictions for multi-dimensional simulations of thermonuclear supernovae generally rely on the tracer particle method to obtain isotopic information of the ejected material for a given supernova simulation.  We investigate how many tracer particles are required to determine converged integrated total nucleosynthetic yields.  For this purpose, we conduct a resolution study in the number of tracer particles for different hydrodynamical explosion models at fixed spatial resolution.  We perform hydrodynamic simulations on a co-expanding Eulerian grid in two dimensions assuming rotational symmetry for both pure deflagration and delayed detonation Type Ia supernova explosions.  Within a given explosion model, we vary the number of tracer particles to determine the minimum needed for the method to give a robust prediction of the integrated yields of the most abundant nuclides.  For the first time, we relax the usual assumption of constant tracer particle mass and introduce a radially varying distribution of tracer particle masses.  We find that the nucleosynthetic yields of the most abundant species (mass fraction $ > 10^{-5}$) are reasonably well predicted for a tracer number as small as 32 per axis and direction -- more or less independent of the explosion model.  We conclude that the number of tracer particles that were used in extant published works appear to have been sufficient as far as integrated yields are concerned for the most copiously produced nuclides. Additionally we find that a suitably chosen tracer mass distribution can improve convergence for nuclei produced in the outer layer of the supernova where the constant tracer mass prescription suffers from poor spatial resolution. ",
        "introduction": "\\label{sec:int} The basic processes of Type~Ia supernovae (SNe~Ia) have been proposed almost 50 years ago \\citep[see e.g.][for a review]{hillebrandt2000a}: A thermonuclear explosion in electron-degenerate matter \\citep{hoyle1960a} produces radioactive \\nuc{56}{Ni} that by its decay delivers energy at exactly the rate observed in SN~Ia light curves \\citep{pankey1962a,truran1967a, colgate1969a, kuchner1994a}.  Despite this long history, the questions of the progenitor system and the explosion scenario are not completely answered.  Pure deflagrations in Chandrasekhar-mass white dwarfs, currently somewhat disfavoured for their apparent inability to produce bright explosions, have been one of the contending explosion scenarios for a long time \\citep[e.g.][]{nomoto1984a,gamezo2003a,roepke2007c}.  If the initial deflagration can transition into a detonation \\citep[e.g.][]{khokhlov1997a,roepke2007d,woosley2007a,woosley2009a}, then better agreement of the models with observations can be obtained \\citep[e.g.][]{roepke2007b,bravo2008a,kasen2009a}.  If such a deflagration to detonation transition does not occur \\citep[cf.\\ e.g.][]{niemeyer1999a}, the possibility that off-center ignition in a single spot leads to a detonation of a white dwarf out of hydrostatic equilibrium in the so-called GCD model is yet another proposed mechanism \\citep{plewa2004a,jordan2008a,meakin2009a}, although the robustness of the initiation of the detonation in this model is also far from certain \\citep[cf.][]{roepke2007a, seitenzahl2009c, seitenzahl2009b}.  However, there are indications from recent stellar population synthesis studies \\citep[e.g.][]{ruiter2009a,mennekens2010a} and X-ray observations of elliptical galaxies and galaxy bulges \\citep{gilfanov2010a}, that the long favoured single degenerate Chandrasekhar-mass progenitor channel is unable to account for observational rates of SNe Ia \\citep{cappellaro1999a}.  Supernovae resulting from the merger of two white dwarfs \\citep[e.g.][]{pakmor2010a} have more favourable statistics and remain a possible explosion channel.  Last but not least, the double detonation sub-Chandrasekhar mass model, which has received renewed interest of late \\citep{fink2007a,fink2010a, sim2010a}, is currently again considered a serious alternative.  The different nucleosynthesis occurring in all these explosion scenarios is an important test for the validity of the respective models.  For one-dimensional simulations, the nucleosynthesis can be calculated during the simulation via the coupling of a nuclear reaction network to the hydrodynamics. Since today's most promising explosion models either explicitly break spherical symmetry (e.g.\\ mergers, double detonations, off-center ignitions) or, in the case of centrally ignited spherically symmetric explosions include buoyancy driven turbulent combustion, at least two-dimensional simulations are required to simulate the essential physics of the explosion. Unfortunately, in highly resolved two-dimensional simulations and especially in three-dimensional simulations use of an extended, full nuclear reaction network during the hydrodynamic evolution is computationally not feasible.  Consequently, the nuclear energy release for multi-dimensional simulations is modeled with simplified and approximate schemes \\citep[e.g.][]{khokhlov1995a,reinecke2002b,vladimirova2006a,calder2007a,townsley2007a,seitenzahl2009a}. Although some works on multi-dimensional explosion simulations of SN Ia do not calculate any detailed nucleosynthesis at all \\citep[e.g.][]{reinecke2002d,plewa2004a,gamezo2005a,plewa2007a,jordan2008a,bravo2008a}, a detailed isotopic composition of the ejecta is often determined in a post-processing step \\citep[][]{travaglio2004a,travaglio2005a,brown2005a,roepke2006b,fink2010a,maeda2010a}. While alternative methods for obtaining isotopic nucleosynthetic yields for SNe Ia are currently under investigation \\citep{seitenzahl2008b,meakin2009a}, the tracer particle method still remains currently the only viable choice for the task.  In Section~\\ref{sec:tpm} we briefly review the tracer particle method and introduce how we distribute tracer particles of variable mass in our initial models.  In Section~\\ref{sec:sim} we give details about our hydrodynamic simulations.  Section~\\ref{sec:conv} presents the results of a convergence study in the number of tracer particles and highlights the advantages of using a radially varying distribution of tracer masses.  We conclude the paper with a summary and discussion in Section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc} We have performed a resolution study in the number of tracer particles required to converge on the integrated nucleosynthetic yield calculations for three different SN Ia explosion models.  We find that for all explosion models, total nucleosynthetic yields appear to have converged (with few exceptions, see Section~\\ref{sec:conv}) for abundant nuclides with mass fractions $> 10^{-5}$ to better than $2\\%$ at 128 tracer particles per axis and direction.  Agreement at better than the $5\\%$ level is already achieved for many of the most abundant nuclides at particle numbers as low as 32 per axis and direction.  Based on these results, we extrapolate that published yields from the literature \\citep[][]{travaglio2004a,travaglio2005a,brown2005a,roepke2006b,maeda2010a} are based on sufficiently high numbers of tracer particles per axis that an adequate to good prediction of the most abundant nuclei is given.  Furthermore, we have shown that for isotopes whose origin is near the surface of the WD, such as e.g. \\nuc{20}{Ne}, \\nuc{26}{Mg} or \\nuc{27}{Al}, better convergence can be achieved if the constraint of constant tracer mass is relaxed and a choice for the tracer particle masses is made that better spatially resolves the outer layers.  The question of how many particles are required to get converged spectra and light curves from subsequent radiative transfer calculations is related.  It seems plausible that a convergence of the total integrated yields is a necessary but not sufficient constraint. For the radiative transfer not only the total mass of a given isotope is important but also its location in velocity space.  It seems that the method of variable mass tracer particles is also promising in this context.  Especially simulations where the nucleosynthesis is dominated by a detonation should greatly benefit from a variable tracer mass distribution, since the mass fractions in the inner regions vary smoothly radially and therefore a de-refinement in the inner part is not a high price to pay for increased resolution in the outer part.  However, the reconstruction of an isotopic abundance for every computational cell of the radiative transfer calculation from the tracer particle distribution is not unique.  The resultant spectra and light curves depend somewhat on particular choice of reconstruction algorithm used. These questions will likely be addressed in a future study."
    },
    "1005/1005.5592_arXiv.txt": {
        "abstract": "We investigate the virialization of cosmic structures in the framework of flat FLRW cosmological models, in which the vacuum energy density evolves with time. In particular, our analysis focuses on the study of spherical matter perturbations, as they decouple from the background expansion, ``turn around'' and finally collapse. We generalize the spherical collapse model in the case when the vacuum energy is a running function of the Hubble rate, $\\CC=\\CC(H)$. A particularly well motivated model of this type is the so-called quantum field vacuum, in which $\\CC(H)$ is a quadratic function, $\\CC(H)=n_0+n_2\\,H^2$, with $n_0\\neq 0$. This model was previously studied by our team using the latest high quality cosmological data to constrain its free parameters, as well as the predicted cluster formation rate. It turns out that the corresponding Hubble expansion history resembles that of the traditional $\\Lambda$CDM cosmology. We use this $\\Lambda(t)$CDM framework to illustrate the fact that the properties of the spherical collapse model (virial density, collapse factor, etc.) depend on the choice of the considered vacuum energy (homogeneous or clustered). In particular, if the distribution of the vacuum energy is clustered, then, under specific conditions, we can produce more concentrated structures with respect to the homogeneous vacuum energy case. ",
        "introduction": "Several cosmological observations (supernovae type Ia, CMB, galaxy clustering, etc.) have converged to a paradigm of a cosmic expansion history that involves a spatially flat geometry and a recently initiated accelerated expansion of the universe (see \\cite{Teg04,Spergel07,essence,Kowal08,Hic09,komatsu08} and references therein). From a theoretical point of view, an easy way to explain such expansion is to consider an additional energy component, usually called dark energy (DE) with negative pressure, that dominates the universe at late times. The simplest DE candidate corresponds to the cosmological constant (see \\cite{Weinberg89,Peebles03,Pad03} for reviews). An elegant model that fits accurately the current observational data, is the spatially flat concordance $\\Lambda$CDM model, which includes cold dark matter (DM) and a cosmological constant, $\\Lambda$. However, the $\\Lambda$ model suffers, among other \\cite{Peri08}, form two fundamental problems: (a) {\\it The ``old'' cosmological constant problem} (or {\\it fine tuning problem}) i.e., the fact that the observed value of the vacuum energy density ($\\rho_{\\Lambda}=\\Lambda c^{2}/8\\pi G\\simeq 10^{-47}\\,GeV^4$) is many orders of magnitude below the value found using quantum field theory \\cite{Weinberg89}, and (b) {\\it the coincidence problem}\\,\\cite{coincidence} i.e., the fact that the matter energy density and the vacuum energy density are of the same order (just prior to the present epoch), despite the fact that the former is a rapidly decreasing function of time while the latter is stationary.  Many authors have attempted to solve the above problems (see \\cite{Peebles03,Pad03,Egan08} and references therein), the key approach being the replacement of the constant vacuum energy either with a time evolving DE (quintessence and the like\\,\\cite{Peebles03}), or alternatively with a time varying vacuum energy density, $\\rho_{\\Lambda}(t)$ \\cite{Shap00,Shap02,Reuter00,Babic02,Grande06,Bas09b}. In the original scalar field models\\,\\cite{Dolgov82} and later in the quintessence context, one can ad-hoc introduce an adjusting or tracker scalar field $\\phi$\\,\\cite{Caldwell98} (different from the usual SM Higgs field), rolling down the potential energy $V(\\phi)$, which could mimic the DE \\cite{Peebles03,Pad03,Jassal,SR,Xin,SVJ}. However, it was realized that the idea of a scalar field rolling down some suitable potential does not really solve the problem because $\\phi$ has to be some high energy field of a Grand Unified Theory (GUT), and this leads to an unnaturally small value of its mass, namely one which is beyond all conceivable standards in Particle Physics. As an example, utilizing the simplest form for the potential of the scalar field, $V(\\phi)=m_{\\phi}^2\\,\\phi{^2}/2$, the present value of the associated vacuum energy density is $\\rho_{\\CC}=\\langle V(\\phi) \\rangle\\sim 10^{-11}\\,eV^4$, so for $\\langle \\phi\\rangle$ of order of a typical GUT scale near the Planck mass, $M_P\\sim 10^{19}$ GeV, the corresponding mass of $\\phi$ is expected in the ballpark of $m_{\\phi}\\sim\\,H_0\\sim 10^{-33}\\,eV$.  Notice that the presence of such a tiny mass scale in scalar field models of DE is generally expected also on the basis of structure formation arguments\\,\\cite{Mota04,Nunes06}; namely from the fact that the DE perturbations % seem to play an insignificant role in structure formation for scales well below the sound horizon. The main reason for this homogeneity of the DE is the flatness of the potential, which is necessary to produce a cosmic acceleration. Being the mass associated to the scalar field fluctuation, proportional to the second derivative of the potential itself, it follows that $m_{\\phi}$ will be very small and one expects that the magnitude of DE fluctuations induced by $\\phi$ should be appreciable only on length scales of the order of the horizon. Thus, equating the spatial scale of these fluctuations to the Compton wavelength of $\\phi$ (hence to the inverse of its mass) it follows once more that $m_{\\phi}\\lesssim\\,H_0\\sim 10^{-33}\\,eV$. All in all, it appears that the problem that one is creating along with the introduction of $\\phi$ is far more worrisome than the problem one is intending to solve, for one is postulating a mass scale which is $30$ orders of magnitude below the mass scale associated to the value of the vacuum energy density ($m_{\\Lambda}\\equiv\\rho_{\\Lambda}^{1/4}\\sim 2.3\\times 10^{-3}\\,eV$).  The analysis of the recent cosmological observations indicates that the DE equation of state (EOS) parameter $w (\\equiv P_{\\rm DE}/\\rho_{\\rm DE})$ is close to $-1$ to within $\\pm 10\\%$, if it is assumed to be constant \\cite{Teg04,Spergel07,essence,Kowal08,Hic09,komatsu08}, whilst it is much more poorly constrained if it varies with time\\,\\cite{essence}. More than two decades ago, Ozer \\& Taha \\cite{Oze87} proposed a time varying $\\Lambda$ as a possible candidate to solve the two fundamental cosmological puzzles, see also \\cite{Free187,Carvalho92} and references therein. In this cosmological paradigm, the dark energy EOS parameter $w$ is strictly equal to -1, but the vacuum energy density (or $\\Lambda$) does evolve with time. Of course, the weak point in this approach % is the unknown functional form of $\\Lambda(t)$, which is however also the case for the vast majority of the DE models. Indeed, in the aforementioned $\\Lambda(t)$ models, the  evolution law is purely phenomenological\\,\\cite{Overduin98}, without a concrete % link to fundamental physics, such as the Quantum Field Theory (QFT) in Curved Space time\\,\\cite{ParkerToms09}. As emphasized in \\cite{ShapSol09}, a completely consistent formulation along these lines should eventually be developed, and such investigations could well be at the heart of one the most important endeavors of theoretical cosmology in the years to come. Therefore, the study of cosmic perturbations in these models is very important\\,\\,\\cite{Grande08,GSFS10} as they might reveal surprises not foreseen in the context of the scalar field models. The new effects may have impact both on the cosmological and the \\textit{astrophysical} domains. While we recently analyzed the potential implications on the former\\,\\cite{Bas09c}, here we focus on the latter domain.  A pioneering QFT fundamental approach to variable $\\CC$ models was actually proposed long ago within the context of the renormalization group (hereafter RG) in\\, \\cite{NelPan82,Toms83}. Later on, the RG-running framework was further explored from different points of view in \\cite{Shap00,Shap02,Reuter00}, and a more systematic presentation from the viewpoint of QFT in curved space-time by employing the standard perturbation RG-techniques of Particle Physics appeared in\\,\\cite{Shap02,Babic02}. Subsequent elaborations, and comparison with the observational tests, confirmed the phenomenological viability of this approach \\,\\cite{RGTypeIa,SSS04,SS12,Fossil07,Bauer05,FSS09a,BFLWard}.  In the class of RG models we shall focus on, the vacuum energy density is expected to vary with time according to the law\\,\\cite{Shap02,RGTypeIa,SSS04,SS12,Fossil07}: $\\CC=n_0+n_2\\,H^2$ (hereafter called the $\\Lambda_{RG}$ model or quantum field vacuum model). As already mentioned, in Ref.\\,\\cite{Bas09c} we have investigated thoroughly the global dynamics of this cosmological model (together with various alternative $\\Lambda(t)$ models), in the light of the most recent cosmological data. However, a serious issue here is how the main features of the largest collapsed cosmic structures, i.e., galaxy clusters, are affected by a running vacuum energy density. We have argued above that this problem can be addressed in scalar field models of the DE, but only at the expense of admitting extremely tiny mass scales which are uncommon in Particle Physics. In this paper, we wish to further explore the alternative option in which the DE component is a time evolving cosmological term $\\Lambda=\\Lambda(t)$, and in this way to assess if the clustering properties of the vacuum energy can shed some light on the fundamental issue of structure formation.  The so called spherical collapse model\\,\\cite{Gunn72}, which has a long history in cosmology, is a simple but still a fundamental tool used to describe the growth of bound systems in the universe via gravitation instability\\,\\cite{Pee93}. In the last decade many authors have studied the small scale dynamics using this model and found that the main features of the cosmic structures (collapse factor, virial density, etc) can potentially be affected by the presence of dark energy \\cite{Lahav91,Wang98,Iliev01,Lokas,Battye03,Maini03,Bas03,Wein03, Mota04,manera,Horel05,Zeng05,Maor05,Perc05,Nunes06,Wang06,david,Fran08, Basi07,Shaef08,Lee09,Abramo}. The aim of the present work is to generalize the spherical collapse model within the variable $\\Lambda_{RG}$ cosmological model, in order to understand non-linear structure formation in such cosmologies and investigate the differences with the respect to the expectations of the concordance $\\Lambda$CDM cosmology.  The structure of the paper is as follows. The basic theoretical elements of the problem are presented in section 2, where we introduce [for a spatially flat Friedmann-Lema\\^\\i tre-Robertson-Walker (FLRW) geometry] the basic cosmological equations. In section 3  we generalize the virial theorem in the case of the QFT $\\Lambda(t)$ cosmological model. Section 4 outlines the theoretical analysis of the spherical collapse model in which the vacuum energy density varies with the cosmic time, and in section 5 we compare the corresponding theoretical predictions of the different models and present a first attempt to use observational data to constrain the different models. We draw our conclusions in section 6. In the appendix A we remind the reader of some basic elements of the concordance $\\Lambda$CDM model in order to appreciate the fact that the $\\Lambda_{RG}$ cosmology is an interesting extension of the standard model. Finally, in appendix B we provide some basic mathematical formulae, while in appendix C we provide accurate fitting formulae for a few important parameters, ie., the density contrast at the turn around redshift and at the epoch of virialization, which do not have a simple fully analytical form. ",
        "conclusions": "In this paper we have studied analytically and numerically the spherical collapse model in the case of a time varying vacuum, with $\\CC(H)=n_0+n_2\\,H^2$, for a spatially flat FLRW geometry. We find that the amplitude and the shape of the virial density contrast is affected by the considered status of the vacuum energy model (homogeneous or clustered). We verify that in the case where the distribution of the vacuum energy is clustered the structures produced are more concentrated (under specific conditions) with respect to the homogeneous dark energy case. Finally, % by comparing the predicted virial density contrast at the present epoch with a preliminary analysis of a suitable subsample of the 2MASS High Density Contrast group catalog (at a mean redshift of $\\langle z \\rangle \\simeq 0.015$), we find that the inhomogeneous vacuum energy models can be accommodated, at a $2\\sigma$ level, if the vacuum clustering parameter is within the range: $-0.009\\mincir \\gamma_{s}\\mincir 0.012$. The latter result points to the direction that perhaps the $\\Delta_{vir}$ parameter, once estimated accurately from observations, could be used in order to determine the internal physical properties of the vacuum energy.  \\vspace {0.5cm} {\\bf Acknowledgments.} JS has been supported in part by MEC and FEDER under project FPA2007-66665, by the Spanish Consolider-Ingenio 2010 program CPAN CSD2007-00042 and by DIUE/CUR Generalitat de Catalunya under project 2009SGR502. MP acknowledges funding by Mexican CONACyT grant 2005-49878.   \\appendix"
    },
    "1005/1005.0717.txt": {
        "abstract": "We undertake a comprehensive and rigorous analytic study of the evolution of radial profiles of covariant scalars in regular Lema\\^\\i tre--Tolman--Bondi dust models. We consider specifically the phenomenon of ``profile inversions'' in which an initial clump profile of density, spatial curvature or the expansion scalar, might evolve into a void profile (and vice versa).  Previous work in the literature on models with density void profiles and/or allowing for density profile inversions is given full generalization, with some erroneous results corrected. We prove rigorously that if an evolution without shell crossings is assumed, then only the `clump to void' inversion can occur in density profiles, and only in hyperbolic models or regions with negative spatial curvature. The profiles of spatial curvature follow similar patterns as those of the density, with `clump to void' inversions only possible for hyperbolic models or regions. However, profiles of the expansion scalar are less restrictive, with profile inversions necessarily taking place in elliptic models.  We also examine radial profiles in special LTB configurations:  closed elliptic models, models with a simultaneous big bang singularity, as well as a locally collapsing elliptic region surrounded by an expanding hyperbolic background. The general analytic statements that we obtain allow for setting up the right initial conditions to construct fully regular LTB models with any specific qualitative requirements for the profiles of all scalars and their time evolution. The results presented can be very useful in guiding future numerical work on these models and in revising previous analytic work on all their applications. ",
        "introduction": "\\footnote{In order to motivate the reading of this long article, we have concentrated most of the background material in the Appendices and have written a proper summary of our results and their implications in the final section (section 11). Readers eager to know these results without going through the technical detail are advised to go directly to this section.} % The well known spherically symmetric LTB dust models \\cite{LTB} are among the most useful known exact solutions of Einstein's equations. There is an extensive literature (see \\cite{kras1,kras2} for comprehensive reviews) using these models, as they allow for probing non--linear effects in inhomogeneous sources by means of analytic and/or tractable numeric solutions. In most applications LTB dust solutions lead to simple but effective toy models of cosmological inhomogeneities, as for example in the series of papers \\cite{KH1,KH2,KH3,KH4,BKH}. The models are also useful in other theoretical contexts, such as censorship of singularities \\cite{sscoll,joshi} and even quantum gravity \\cite{quantum}. More recently, LTB models have been extensively used in the widespread effort to explain cosmic dynamics without resorting to an elusive source like dark energy, either by fitting observations \\cite{LTB1,LTB2,LTBkolb,num1,num2}, or in the context of scalar averaging of inhomogeneities \\cite{LTBfin,LTBchin,LTBave1,LTBave2,LTBave3} (see \\cite{celerier,ave_review} for comprehensive reviews on these applications). For a novel theoretical approach to the dynamics of these models see \\cite{wainwright}.  Since large scale cosmic structure is dominated by large voids, the study and understanding of the spatial profiles of density and velocity in inhomogeneities is a relevant theoretical and practical issue that could provide important clues on the formation of these voids as part of the dynamics of cosmic sources. While LTB models provide an idealized description of cosmic inhomogeneities, a proper understanding (even qualitative) of radial profiles of scalars in these models still has a significant potential in astrophysical and cosmological implications in this context \\cite{kras1,kras2} (see \\cite{CBK} for an update). In particular, knowing how these radial profiles evolve is extremely useful when applying to these models a formalism of scalar averaging of inhomogeneities, such as Buchert's formalism \\cite{LTBave2,LTBave3}.  An important theoretical issue in looking at density radial profiles is the possibility that a concavity change could occur under regular conditions (particularly around a symmetry center). This implies that a given initial density profile with a ``clump'' (over--density) form could evolve into a ``void'' (under--density) profile, or vice versa, a phenomenon that is usually denoted by a ``clump to void'' or ``void to clump'' density profile inversion. In a well known article Mustapha and Hellaby \\cite{mushel} claimed to have furnished proof for the existence of this profile inversion, and to have found the conditions in which it occurs. However, since the conditions provided by these authors are too complicated (involving second and third order radial derivatives), their proof (their section 5) is based on restrictive assumptions, and thus is not applicable to generic LTB models. In the end, they discuss the observational implications of these profile inversion and furnish basically an empiric proof of their existence by looking at specific models and numeric examples where it occurs.  Following a completely different methodology from that of Mustapha and Hellaby, we correct, extend and fully generalize their work in this article. By considering fully generic and regular LTB models, we examine the evolution of the radial profiles of the basic covariant scalars: rest mass density, $\\rho$, spatial curvature, $\\RR$, and Hubble expansion scalar, $\\Theta$, aiming  specifically at dealing in full generality with the following issues:\\, (i) to provide a precise characterization of radial profiles as ``clumps'' or ``voids'', and (ii) to study in full rigor the conditions that allow for the existence of profile inversions (assuming full regularity). As the reader can find out by going through the article (or by looking at the summary of results in section 11), we do find fully analytic and rigorous conditions for the existence of clump/void radial profiles and their possible inversions for regular LTB models in general.  An interesting and similar approach to the study of the evolution of radial profiles is found in the articles by Krasinski and Hellaby \\cite{KH1,KH4} (see the review in \\cite{kras2}). These authors find the conditions for the existence of a unique LTB model, such that a given density or ``velocity'' profile at a given time constant hypersurface can be ``mapped'' to any other profile at a second hypersurface. While there is an obvious intersection and complementarity between the results of this article and those of \\cite{KH1,KH4}, our methodology is different and (in our opinion) less restrictive, as we only require generic initial data to be specified at a single time constant hypersurface that can be considered a Cauchy surface in the context of an initial value problem.  A summary of articles that have examined LTB models with void profiles in a perturbative and non--perturbative context can be found in the extensive review of this literature (previous to 1994) given in section 3.1 of \\cite{kras1}, and in the update in section 18.5 of \\cite{kras2}. Some authors \\cite{occhio81,occhio83} have examined non--perturbative void profiles numerically in the context of structure formation scenarios (though, from our results in section 6, the elliptic models in \\cite{occhio81} cannot be free from shell crossings). In some suggested configurations the void is modeled as an under--dense FLRW region, which is connected to a cosmic ``background'' through a section of a closed elliptic model \\cite{cham,sato}. In other articles \\cite{boncham1,boncham2,meszaros}, the void is an LTB region containing a center, matched at a fixed comoving radius to another LTB region (the ``envelope''), which in turn is matched to a FLRW ``exterior''. Although these models are rather artificial, their results agree with those of this article.  The contents of the article is summarized in the remaining of this section. Section 2 contains the bare basic background material needed to understand the article. The known analytic solutions of the field equations (given in Appendix A), which are determined by a preferred set of free functions ($M,\\,E,\\,\\tbb$), have become the standard formulation in practically all theoretical, empiric and in numeric work on these models. While this standard parametrization is adequate and works in practice, we utilize an alternative set of covariant quasi--local scalars that are more suitable for our purpose (and that can be very convenient for numerical work, as well as a better understanding of several theoretical and practical issues)  \\cite{sussQL1,sussQL,suss08,suss09,suss10a,suss10b}. We also express these scalars and their fluctuations as quantities that are scaled with respect to their initial values in the framework of an initial value problem \\cite{suss10a}.  Extra background material, which is important as well but can be distracting if it appears in the main body of the article, has been placed in Appendices B, C and D. For example, the parametrization of the known analytic solutions in terms of the new variables and its initial value formulation (Appendix B), the Hellaby--Lake conditions \\cite{suss10a,HLconds,ltbstuff} that guarantee an evolution free from shell crossings (Appendix C), and a discussion (Appendix D) of the relation between the radial coordinate and the proper radial length, which is the affine parameter along radial rays \\cite{suss10b}. This is important since we are looking at the behavior of covariant scalars in the radial direction.  We provide in section 3 a precise and rigorous characterization of the ``clump'' or ``void'' nature of a radial profile of a scalar $A$ in terms of the sign of the  gradient $A'$ (monotonicity of $A$), which leads to the definition of a ``turning value'' (to be denoted ``TV of $A$'') as the value $r=\\rtv$ where this gradient vanishes. In section 4 we examine the compatibility between initial clump/void profiles and the Hellaby--Lake conditions, while the concept of ``profile inversion'' is defined in section 5 in terms of the vanishing of the fluctuations defined in section 2, something which can occur with or without the existence of TV's.  In section 6 we examine in detail density radial profiles and provide necessary and sufficient conditions for density profile inversions for parabolic, hyperbolic and open elliptic models or regions. We remark that the locus where the density fluctuation vanishes when there are density TV's corresponds to the ``density wave'' (see section 3.1 of \\cite{kras1}). The profiles and profile inversions for spatial curvature and the expansion scalar are respectively examined in sections 7 and 8. The case of closed elliptic models is considered in section 9, as for these models the spherical topology of the hypersurfaces of constant time introduces a TV of the area distance $R$. In section 10 we look at radial profiles and profile inversions in special LTB configurations: simultaneous big--bang singularity, LTB models or regions not containing symmetry centers and configurations constructed by glueing LTB regions with ``mixed'' kinematics (elliptic, parabolic and hyperbolic regions, see \\cite{ltbstuff}), and specially the mixed configuration made by an elliptic ``interior'' region surrounded by a hyperbolic ``exterior''.  Section 11 provides a summary and final discussion.  Appendix E describes an alternative parametrization of the models in terms of quantities that can be related to the Omega and a Hubble parameters of a FLRW cosmology in the homogeneous limit. It is straightforward to translate the results of this article in terms of these parameters, which have been used in various articles \\cite{num1,LTBfin,suss08} describing cosmological inhomogeneities in various contexts (see also \\cite{suss10a}). Appendix F provides the explicit expressions for the basic covariant scalars that characterize LTB models. ",
        "conclusions": "We have conducted a comprehensive and rigorous examination of the radial profiles of the main covariant scalars (density, spatial curvature and expansion, respectively denoted by $m,\\,k,\\,\\HH$, see definition (\\ref{mkHdefs})) for regular LTB models in full generality. The time evolution of these profiles has been considered, and in particular, we have addressed the issue of profile inversions in which an initial clump can evolve into a void (or vice versa). The necessary background material is the formalism of quasi--local variables within an initial value formulation (section 3), rephrasing in terms of this formalism the analytic solutions (Appendix B) and the Hellaby--Lake conditions (Appendix C), and the relation between the radial coordinate and proper radial length (Appendix D). The quasi--local density ($m_q$), spatial curvature ($k_q$) and expansion scalar ($\\HH_q$) are given by \\ref{mq})--(\\ref{Hq}), while the fluctuations $\\Da$ follow from (\\ref{Dadef}).  \\subsection{Summary of results.}  We have used thoughout the article the term ``turning value'' of any scalar $A$ (``TV of $A$'') to denote a value $r=\\rtv$ such that $A'$ vanishes at $r=\\rtv$ at a given hypersurface $\\T[t]$ marked by constant $t$ (this definition also applies to the $A_q$). In general, the value $\\rtv$ is different for each $\\T[t]$, but it can be the same for all $t$ (as for example a regular TV of $R$). Throughout sections 3--8 we have considered only LTB models or regions containing a symmetry center and with an open topology ($R'>0$ holds everywhere: no TV of $R$), leaving the case with closed topology (regular TV of $R$) and special configurations for sections 9 and 10. The main results contained in 10 lemmas that were proven in these sections are summarized below:  \\begin{itemize}  \\item {\\bf{Section 3}}. The clump/void character of radial profiles of local and quasi--local scalars, $A$ and $A_q$, in radial domains containing a center was given in terms of the existence of turning values, TV's. Assuming that $R'>0$ holds, we proved the following lemmas: % \\begin{itemize} \\item {\\bf{Lemma 1}}: The signs of the radial gradient $A'_q$ is related to the sign of the fluctuation $\\Da$. \\item {\\bf{Lemma 2}}: The necessary and sufficient condition for a TV of $A_q$ at some $r=\\rtv$ is given by $\\Da=0$ at $r=\\rtv$. In general, $\\rtv=\\rtv[t]$ is different at different times (see figure 1). \\item {\\bf{Lemma 3}}: The existence of a TV of $A_q$ in a radial domain is a sufficient condition for the existence of a TV of $A$ in the same domain. \\end{itemize} % These lemmas allow us to examine the monotonicity of scalars by looking at signs and the zeros of their fluctuations. As a consequence of Lemma 3, it is sufficient to examine the profiles of the scalars $A_q$ (which satisfy simpler scaling laws) to know the profiles of the scalars $A$.  \\item {\\bf{Section 4}}. Character of initial density and spatial curvature profiles compatible with the Hellaby--Lake conditions. Initial density must have a clump profile in parabolic and elliptic models or regions, while initial clump and void profiles are possible in hyperbolic models or regions. Initial curvature  must have a clump profile in elliptic models or regions, but hyperbolic models or regions admit initial clumps and voids.  Explicit conditions are given for each case.  \\item {\\bf{Section 5}}. Formal definition of profile inversion and its relation with TV's of scalars. Profile inversions can occur with and without TV's. The following lemmas were proven: % \\begin{itemize} \\item {\\bf{Lemma 4}}: A profile inversion of $A_q$ implies a profile inversion of $A$. \\item {\\bf{Lemma 5}}: Sufficient conditions for the existence of a profile inversion of $A_q$ follow from the existence of TVs of $A_q$ plus some extra requirements on the fluctuations $\\Da$. This type of profile inversion is illustrated by figure \\ref{fig1}. \\item {\\bf{Lemma 6}}: Sufficient conditions for the existence of a profile inversion of $A_q$ when there are no TV's of $A_q$. This type of profile inversion is illustrated by figure \\ref{fig2}. \\end{itemize} % Notice that Lemma 4 guarantees that the results of Lemmas 5 and 6 apply also to the local scalars $A$. The conditions in Lemmas 5 and 6 are independent of the Hellaby--Lake conditions, hence the latter place extra constraints when applying these lemmas to specific configurations. % \\item {\\bf{Section 6}}. Density radial profiles and profile inversions are examined in detail. The following lemmas were proven: % \\begin{itemize} \\item {\\bf{Lemma 7}}: Necessary and sufficient condition are given for the `clump to void' density profile inversion with density TV's in regular hyperbolic models or regions. Void to clump inversion is not possible. \\item {\\bf{Lemma 8}}: Necessary and sufficient condition for the `clump to void' inversion without TV's in regular hyperbolic models or regions. Void to clump inversion is not possible. \\item {\\bf{Lemma 9}}: Density profile inversions, with or without density TV's, cannot occur in regular open elliptic models or regions containing a center (no TV's of $R$). \\end{itemize} % The main result of this section (and possible of the whole article) follows as a consequence of these lemmas:  \\begin{quote} \\noindent {\\it {the only density profile inversion (with and without TV's) compatible with absence of shell--crossings (the Hellaby--Lake conditions) is the `clump to void' inversion in hyperbolic models or regions.}} \\end{quote}  \\noindent The specific conditions outlined by Lemmas 7 and 8 (equations (\\ref{dm32dk}) and (\\ref{CintoV2})), which must be satisfied by hyperbolic models and region, are simple restrictions on their initial value functions (density, spatial curvature and their fluctuations). It is evident that these are not weird or outlandish conditions, but reasonable and easy to prescribe. % \\item {\\bf{Section 7}}. Spatial curvature profiles and their inversions. The only inversion compatible with the Hellaby--Lake conditions is the `clump to void' inversion in hyperbolic models or regions. These inversions only occur in the particular parameter cases when there is no inversion of the density profile. % \\item {\\bf{Table 1}} provides the combination of initial conditions that allow for profile inversions of density and spatial curvature in hyperbolic models or regions. % \\item {\\bf{Section 8}}. Radial profiles and profile inversions of the expansion scalar. Hyperbolic models and regions allow for a regular evolution with and without profile inversions (conditions are provided for each case). However, profile inversions of the expansion scalar necessarily occur in elliptic models or regions containing a center. This is illustrated by figure \\ref{fig3}.  \\item {\\bf{Section 9}}. Closed elliptic models. A TV of $R$ occurs under regular conditions at the same value $r=\\rtv$ for all times. We proved the following lemma: % \\begin{itemize} \\item {\\bf{Lemma 10}}: The TV of $R$ at $r=\\rtv$ implies that $\\rtv$ is a common TV of the local and quasi--local density, spatial curvature and expansion scalar. The converse is not true, as we proved in Lemmas 1--5 and 7--8 that TV's of these scalars can occur when there is no TV of $R$ ($R'>0$ holds everywere). \\end{itemize} % As an important corollary of this lemma, the fluctuations $\\Da$ do not change sign because of the TV of $R$. As a consequence, there are no profile inversions of density and spatial curvature, though profile inversions of the expansion scalar must occur as in open elliptic models. % \\item {\\bf{Section 10}}. Special configurations: % \\begin{itemize} \\item {\\underline{Regular hyperbolic and elliptic models with a simultaneous big bang}}. Profile inversions of density and spatial curvature (with or without TV's) cannot occur. Hyperbolic models allow for a void profile for all the time evolution. % \\item {\\underline{Regions not containing symmetry centers}}. The results of all proven lemmas hold, the only caveat being the fact that (in general) we have $\\Da\\ne 0$ at the boundaries of the radial range (as opposed to $\\Da= 0$ strictly holding at the symmetry center). % \\item {\\underline{LTB models with ``mixed'' kinematics}}. These configurations are made by glueing or matching combinations of parabolic, hyperbolic and elliptic regions (see \\cite{boncham1,boncham2,meszaros,ltbstuff}). The results of all proven lemmas hold. % \\item {\\underline{Elliptic region surrounded by a hyperbolic exterior}}. From Lemma 9 and assuming that Hellaby--Lake conditions hold, the elliptic region must have clump profiles for the density and spatial curvature, without inversions. However, void profiles and `clump to void' inversions are possible in the hyperbolic exterior in agreement with Lemma 7 and the combination of initial conditions specified in Table 1. \\end{itemize}  \\end{itemize}  \\subsection{Relaxation of regularity conditions.}  In obtaining the results summarized above we have assumed that the Hellaby--Lake conditions \\cite{HLconds,ltbstuff,suss10a} (see Appendix C) hold for all the evolution times of the models (shell crossings are completely absent). Evidently, if we relax these conditions by demanding that they hold only for all $t>t_i$ for a given $t=t_i$, then some of the results of the Lemmas are no longer binding. The most important example is Lemma 9, which forbids density void profiles and `clump to void' density profile inversion in elliptic models. This result follows from the fact that a density void profile is incompatible with the joint fulfillment of two of the Hellaby--Lake conditions (\\ref{noshxGe}):\\,  $\\tbb'\\leq 0$ and $\\tcoll'\\geq 0$, where $\\tbb$ and $\\tcoll$ denote the locus of the initial and collapsing singularity. However, if we demand $\\tcoll'\\geq 0$ to hold but not $\\tbb'\\leq 0$, then a shell crossing singularity necessarily emerges for $t\\approx \\tbb$ for all $r$, but not for $t\\approx \\tcoll$. Depending on the free parameters $m_{qi},\\,k_{qi}$, it may be possible to construct an elliptic model that is free from shell crossings for all $t>t_i$ for some $t_i>\\tbb$ \\cite{suss02}, and that admits a `clump to void' inversion and density void profiles in this time range. Since a dust source does not provide a good description of the physical conditions near an initial singularity in a cosmological model, it may be sufficient for a physically reasonable evolution to demand absence of shell crossings for all $t>t_i$, provided one can justify that the dust description beaks down for $t<t_i$.  The relaxation of the Hellaby--Lake conditions, as described above, can also work for hyperbolic models or for special configurations. It can also allow for a density `void to clump' profile inversion, and it can also be set up so that an evolution free from shell crossings occurs in time ranges restricted to an early stage $\\tbb < t<t_i$. This latter case can be useful for studying void formation in the context of toy models of structure formation (``spherical collapse model''), in which the collapsing stage of an over--density with positive spatial curvature (elliptic) is omitted,  as the late stage of the dynamics is dominated by the virialization process leading to the formation of stable structures (hence  there is no collapse stage because the assumption of spherical dust is no longer valid)\\cite{padma}.  \\subsection{Final discussion.}  The most important results in this article are those obtained in section 6 (Lemmas 7, 8 and 9), section 9 (Lemma 10 and its corollaries) and the elliptic/hyperbolic mixed configuration in section 10. These results correct and provide full generality to the work initiated by Mustapha and Hellaby \\cite{mushel}, as these authors only furnished a restricted proof (together with particular and numeric examples) of the existence of the `clump to void' density profile inversion in hyperbolic models, without attempting to prove it for elliptic or parabolic models. They also claimed (mistakenly) that the `void to clump' inversion was possible without shell crossings. Since they only examined density profiles, the results of sections 7 and 8 also extend and generalize their work to spatial curvature and the expansion scalar.  The fact that the existence of density void profiles (with total absence of shell crossings) is only possible in hyperbolic models and regions is consistent with the conclusions of \\cite{boncham1,boncham2,meszaros}. However, as opposed to these authors, we have obtained the analytic conditions for these profiles and also for their inversions in full generality, and without resorting to artificial configurations made by matchings regions in which the free parameters take special forms and radial gradients and density are likely discontinuous.  Since the results of Lemma 9 are analytic and general, the elliptic models with void profiles examined in \\cite{occhio81} must necessarily have shell crossings (which was not proven otherwise by the authors). Also, some of the configurations in the first and second column of Table 1 of \\cite{meszaros} are listed as free from shell crossings. This is mistaken, as the Hellaby--Lake conditions are incompatible with elliptic and parabolic regions that contain a symmetry center and have density void profiles.  Radial profiles of scalars in closed elliptic models are strongly constrained by the TV of $R$ that occurs at a fixed value $r=\\rtv$, which (as proven by Lemma 10) is a common TV for the density, spatial curvature, expansion scalar and all initial value functions (but not the initial fluctuations). Since this TV implies a change of sign of the gradients $m'_q$ and $m'$ at $r=\\rtv$, it introduces relative under--dense or over--dense regions in thick spherical shells of dust layers, as opposed to a compact spherical region around the center. These shells can be considered as some sort of  ``topological'' density clumps or voids (the same remark applies to spatial curvature or the expansion scalar), and are a product of the spherical $\\mathbb{S}^3$ topology of the rest frames $\\T[t]$, and as such are not related to a profile inversion produced by the evolution of the models given generic initial conditions (as is the case for the TV's that were examined in sections 3--8). As a consequence, models of topological clumps/void profiles \\cite{cham,sato} are rather artificial constructions.  It is important to emphasize that the radial profile of the expansion scalar, as examined in section 8, is much less restrictive to the existence of TV's and profile inversions than the profiles of density and spatial curvature. This fact is crucial to assess the existence and magnitude of a ``back--reaction'' term and an ``effective'' acceleration in the context of Buchert's scalar averaging applied to LTB models \\cite{LTBave2,LTBave3}, which is a relevant issue in current research involving LTB models. Also, as pointed out in recent literature \\cite{wiltshire}, spatial gradients of the expansion scalar (related to gradients of the binding energy) are the main generators of back--reaction (see also \\cite{LTBave2,LTBave3}), and as such are important in providing a theoretical interpretation of cosmological observations that avoid introducing a dark energy source.  Regarding the elliptic region surrounded by a hyperbolic external region, the fact that the latter may admit a density void profile (and a `clump to void' inversion), without violating regularity conditions (Hellaby--Lake), is indeed a nice and unexpected result. This possibility leads to simple models in which a high density region undergoing collapse is surrounded by a large spherical shell in an intermediary transition scale, whose density is lower than that of the cosmic large scale asymptotic background. Configurations of this type could be compatible with the notion of ``finite infinity'' in which cosmic bound structures are approximately regarded as asymptotically flat at such intermediate scales between the cosmic voids and the far cosmic background \\cite{wiltshire,fiEllis}.  Finally, it is necessary to remark that we have conducted a theoretical study of radial profiles and profile inversions based on general analytic expressiond (avoiding the use of special cases based on narrow parameter specializations). Hence, we have not taken into consideration the question of whether these radial profiles are compatible or not with the constraints that follow from actual cosmological and astrophysical observations. However, introducing specific profiles meeting these observational criteria is an important task that certainly requires a follow up article based on a full numeric approach, preferably using the Omega and Hubble parameters discussed in Appendix E. This article is presently under elaboration.   \\begin{appendix}"
    },
    "1005/1005.1652_arXiv.txt": {
        "abstract": "We present multi-wavelengths observations and a radiative transfer model of a newly discovered massive circumstellar disk of gas and dust which is one of the largest disks known today. Seen almost edge-on, the disk is resolved in high-resolution near-infrared (NIR) images and appears as a dark lane of high opacity intersecting a bipolar reflection nebula. Based on molecular line observations we estimate the distance to the object to be 3.5 kpc.  This leads to a size for the dark lane of $\\sim$10500 AU but due to shadowing effects the true disk size could be smaller. In \\emph{Spitzer/IRAC} 3.6 $\\mu$m images the elongated shape of the bipolar reflection nebula is still preserved and the bulk of the flux seems to come from disk regions that can be detected due to the slight inclination of the disk. At longer \\emph{IRAC} wavelengths, the flux is mainly coming from the central regions penetrating directly through the dust lane. Interferometric observations of the dust continuum emission at millimeter wavelengths with the SMA confirm this finding as the peak of the unresolved mm-emission coincides perfectly with the peak of the  \\emph{Spitzer/IRAC} 5.8 $\\mu$m flux and the center of the dark lane seen in the NIR images. Simultaneously acquired CO data reveal a molecular outflow along the northern part of the reflection nebula which seems to be the outflow cavity.  An elongated gaseous disk component is also detected and shows signs of rotation. The emission is perpendicular to the molecular outflow and thus parallel to but even more extended than the dark lane in the NIR images. Based on the dust continuum and the CO observations we estimate a disk mass of up to a few solar masses depending on the underlying assumptions. Whether the disk-like structure is an actual accretion disk or rather a larger-scale flattened envelope or pseudodisk is difficult to discriminate with the current dataset. The existence of HCO$^+$/H$^{13}$CO$^+$ emission proofs the presence of dense gas in the disk and the molecules' abundances are similar to those found in other circumstellar disks. We furthermore detected C$_2$H towards the objects and discuss this finding in the context of star formation. Finally, we have performed radiative transfer modeling of the K band scattered light image varying a disk plus outflow 2D density profile and the stellar properties. The model approximately reproduces extent and location of the dark lane, and the basic appearance of the outflow. We discuss our findings in the context of circumstellar disks across all mass regimes and conclude that our discovery is an ideal laboratory to study the early phases in the evolution of massive circumstellar disks surrounding young stellar objects. ",
        "introduction": "In recent years, significant progress has been made in our understanding of Young Stellar Objects (YSOs) and the formation of stars. Still, there remain unanswered questions and problems that have to be solved observationally and theoretically. In particular, while the formation of low-mass and intermediate mass stars by accretion of material via a circumstellar disk is well established \\citep[e.g.,][]{watson2007}, for massive stars (M $\\ge$ 8 M$_\\sun$) not only the theory of the star formation process \\citep[e.g.,][]{stahler2000} but also the observations are challenging \\citep[e.g.,][]{beuther2007,zinnecker2007} and yield partly contradicting results \\citep[e.g.,][]{chini2004,sako2005,steinacker2006,nurnberger2007}. Conceptually the problem arises from the fact that at least in the spherically  symmetric case the strong radiation pressure of a hydrogen burning young massive star with a mass  $\\ge$ 8 M$_\\sun$ would prevent any further gas infall and would thus halt the formation of a more massive object. Two competing theories overcoming this problem were intensively discussed in the recent past. They explain the formation of massive stars either by the so-called \"Competitive Accretion and Merging Model\" \\citep{bonnell2006} or by fragmentation during a large scale collapse in turbulent cores \\citep{mckee2003}. The latter model resembles the formation process of low-mass objects and consequently relies on circumstellar disks around young massive stars to build up the final objects \\citep[see,\t also,][]{krumholz2009}.  Today, there is ample evidence for the existence of \"disk-like\" structures around massive young stellar objects with masses/sizes being higher/bigger compared to typical low-mass circumstellar disks \\citep[e.g.,][]{cesaroni2007,schreyer2006,chini2004,beltran2004,patel2005,jiang2005,jiang2008}. While for early B-stars these structures appear to be indeed rather disk-like, for O-stars they appear to have more of a toroidal shape. In most cases large, flattened structures were detected using high spatial resolution molecular line data in the (sub-)mm regime. Typically, at shorter wavelengths, the dense envelope surrounding the massive protostar causes too much extinction and prevents a direct detection of the source in the mid- or near-infrared. Indications for rotation in the flattened gaseous structures supported the hypothesis that these objects are massive circumstellar disks. However, the question whether these objects are indeed the high-mass counterpart to the well-established low-mass protoplanetary accretion disks directly feeding the central source or whether they are (remnants of) the rotating and infalling molecular envelopes \\citep[e.g.,][]{beuther2009} or \"pseudodisks\"\\footnote{A \"pseudodisk\" is a flattened disequilibrium structure created during the collapse of a magnetized cloud core by strong magnetic pinching forces deflecting the infalling gas toward the equatorial plane \\citep{gallishu1993}.} \\citep{gallishu1993,allen2003} is unclear. In particular, it appears likely that at least some of the observed massive gaseous \"toroids\" around stars of 10-20 M$_\\sun$ are unstable and have comparatively short lifetimes \\citep{cesaroni2007} while low-mass disks typically persist for several Myr.  In addition to (sub-)mm data, only for very few massive objects mid- or near-infrared (MIR / NIR) data exist that support the existence of an extended dusty disk component. \\citet{chini2004} presented an NIR image of an assumed accretion disk surrounding a presumably massive protostar in M17, but their interpretation was challenged by \\citet{sako2005}. A detailed radiative transfer modeling of the $K$-band image of this source was performed in Steinacker et al. (2006) showing that the underlying structure has a disk shape with a mass around a few M$_\\odot$. Using NIR polarimetric images \\citet{jiang2005,jiang2008} found elongated polarized signatures around massive young stellar objects suggesting the presence of polarization disks. \\citet{sridharan2005} claimed the detection of a silhouette edge-on disk around IRAS 20126+4104 based on NIR images in the $K$, $L'$, and $M'$ filters. More recently, however, \\citet{debuizer2007} derived from MIR images at 12.5 and 18.3\\,$\\mu$m a much more complex structure of the inner regions of the same source and suggested that the data could also be interpreted having a tight young stellar cluster in the center. Based on MIR images, \\citet{chini2006} reported the detection of a remnant disk around a young massive star, and \\citet{nielbock2007} reported on evidence for a disk candidate surrounding a hypercompact HII region in M17 based on NIR data. These authors also provided a radiative transfer model for the disk seen in the $K_s$-band image. Finally, very recently \\citet{okamoto2009} published the detection of a flared circumstellar disk around the 10 M$_\\sun$ Herbig Be star HD200776 also using MIR imaging data.  Here, we report the direct detection of a large and massive circumstellar disk of dust and gas surrounding a young stellar source. The object is named CAHA J23056+6016 as it was first identified in NIR images taken at the Calar Alto observatory (CAHA). Currently, we can not say whether the disk is an actual accretion disk or rather a \"pseudodisk\" or a flattened envelope and we use the more general term \"circumstellar disk\" or simply \"disk\" throughout the paper to avoid any possible misinterpretation. The morphology shows a bipolar nebula intersected by a dark lane. Multi-wavelengths data including NIR/MIR images as well as dust continuum and molecular line data are used to constrain the physical properties of this object. We derive a distance of 3.5 kpc to the source, find a radius for the disk of several thousand AU and a mass, depending on the underlying assumptions, of up to several solar masses. The $K-$band and the mm image are analyzed with a disk/outflow model using radiative transfer calculations to derive scattered light images.  In section 2 we describe the observations and data reduction and present the results of the data in section 3. Results from the radiative transfer model are presented in section 4. A discussion about the nature and properties of CAHA J23056+6016 in the context of star formation is given in section 5. Finally, we summarize our findings and conclude in section 6. ",
        "conclusions": "In this paper we presented a new large and massive circumstellar disk (named CAHA J23056+6016) presumably surrounding a young stellar object. Our findings can be summarized as follows: \\begin{itemize} \\item The disk is seen almost edge-on and was detected in high resolution NIR images where it caused a dark lane intersecting a bipolar nebula that is most likely illuminated by a young stellar source. The dark lane stretches over $\\sim 3''$ which, for an assumed distance of 3.5 kpc, translates into $\\sim$10500 AU, but due to shadowing effects the true physical disk size is probably smaller (see also last point). \\item  CAHA J23056+6016 appears to be one of the largest circumstellar disks known today but at the moment we have no empirical evidence whether it is an actual accretion disk or a \"pseudodisk\" resulting from the collapse of a moderatly magnetized rotating molecular cloud core or a flattened remnant envelope. \\item The young stellar source residing behind the disk is not detected in the NIR images as the optical depth through the disk is too high. \\item In \\emph{Spitzer IRAC} images the object is also detected and at 5.8 $\\mu$m one starts probing directly through the disk and detects flux mainly coming from the central disk regions. \\item This finding is supported by the fact that the peak of the emission seen in different \\emph{Spitzer IRAC} filters is shifting with wavelength and that at the longest available \\emph{IRAC} wavelength (5.8 $\\mu$m) the emission peak coincides with the center of the dark lane seen in the NIR images. \\item In interferometric mm dust continuum observations with the SMA we detected but not resolved CAHA J23056+6016 at the 11.1 mJy/beam level. This sets the upper limit for the disk radius to $\\le$\\,7000 AU. \\item The mm dust emission peaks again nicely on the dark lane seen in the NIR images and coincides perfectly with the  \\emph{IRAC} 5.8 $\\mu$m emission. From the unresolved continuum data we estimate the disk mass to be in the order of 1 M$_\\sun$ making it one of the rare cases where a massive disk is resolved in NIR images. \\item We found a molecular outflow in CO(2--1) coinciding perfectly with the northern part of the bipolar nebula seen in the NIR images. In addition, we found indications for a gaseous disk component in C$^{18}$O data, with the emission being perpendicular to the outflow and showing signs for rotation. The gaseous disk is even larger than that derived from the NIR images or the dust continuum observations and yields a mass of $\\sim$4.5 M$_\\sun$. \\item We also detected H$^{13}$CO$^+$ and HCO$^+$ line emission from CAHA J23056+6016 which arises from regions of higher density. The derived abundances are comparable to those found in other circumstellar disks but lower compared to protostellar source. \\item The detection of C$_2$H supports a young age for the object. We suggest that the emission could either arise from the disk itself but could also be emitted from externally illuminated regions of the disk / envelope system that are subject to high UV radiation fields from nearby HII regions. \\item The current data does not allow us to constrain the mass of the central object directly but, based of the derived disk mass, we speculate that it could have up to several solar masses as well. \\item Interestingly, we find no signs of a young cluster around CAHA J23056+6016 within the typical cluster size of 0.5\\,--\\,1 pc. If the central source were rather massive as well then we had found a rare case of isolated massive star formation. \\item Using a simple 2D dust density profile for the disk and the reflection nebula, the radiative transfer modeling of the $K$-band image could constrain the inclination of the disk to be 5\\,--\\,15$^\\circ$ from an edge-on view. The scale height ratio of the disk ranges from 0.1 to 0.3 to be consistent with the extent of the dark lane. The radial power-law index was not as well constrained as the inner and outer radius of the disk, where the latter parameter is found to be $\\sim$4000 AU. \\end{itemize} The disk of CAHA J23056+6016 is extremely large and very massive and the fact that we have spatially resolved information at a variety of wavelengths makes this object an ideal laboratory to study the properties of a disk surrounding a young presumably intermediate mass or even massive star. Our finding argues for a star formation process that is functioning over a wide mass range, ranging from sub-stellar objects \\citep[e.g.,][]{zapatero2007,quanz2010} to, as it appears, at least several solar masses. In addition, as higher mass stars seem to process their circumstellar disks faster than low-mass objects \\citep[e.g.,][]{alonso-albi2009}, CAHA J23056+6016 provides the opportunity to study an early stage of disk evolution where the disk is still massive. %  For future observations of CAHA J23056+6016 it might be interesting to aim at an even higher spatial resolution of the dust continuum emission to eventually resolve the disk. Filling in the SED in sub-mm and MIR wavelengths could in addition help to even better characterize the disk geometry and dust properties. In parallel, observations of the gaseous disk component with higher spectral resolution will shed additional light on the detected rotation signatures. Finally, it might be worth to obtain a NIR spectrum of the reflection nebula. This might allow us to search for spectral features that constrain the spectral type of the central source and possible bear signs of accretion processes."
    },
    "1005/1005.4288_arXiv.txt": {
        "abstract": "We propose an estimator defined in real space for the reconstruction of the weak lensing potential due to the intervening large scale structure from high resolution maps of the cosmic microwave background. This estimator was motivated as an alternative to the quadratic estimator in harmonic space to surpass the difficulties of the analysis of maps containing galactic cuts and point source excisions. Using maps synthesised by pixel remapping, we implement the estimator for two experiments, namely one in the absence and one in the presence of detector noise, and compare the reconstruction of the convergence field with that obtained with the quadratic estimator defined in harmonic space. We find good agreement between the input and the reconstructed power spectra using the proposed real space estimator. We discuss interesting features of the real space estimator and future extensions of this work. ",
        "introduction": "Observations of the cosmic microwave background (CMB) are a powerful probe of the physics of the early universe and a robust discriminant of cosmological models. Since the CMB photons propagate more or less freely after recombination, what we observe on large scales are mainly the perturbations on the last-scattering surface. The CMB contains also an imprint of the photon quadrupole which developed during the recombination epoch and generated linear polarisation through the mechanism of Thomson scattering.  However, the CMB should also contain signatures of processes that occurred between the last-scattering surface and the present. These processes will be manifest as non-linear signals which become important on small scales. The growing improvement in resolution and sensitivity of CMB experiments allows us to resolve ever so smaller scales and motivates the development of estimators % and % robust implementations for the extraction of the different signals. Among such processes is weak gravitational lensing, which consists of the deflection of CMB photons % by the mass of matter clustered along the line of sight, while leaving the surface brightness unaffected. This effect distorts the temperature-temperature (TT) correlation power spectrum by a few arcminutes but coherently over degree scales \\cite{seljak}, % which means that lensing becomes an important effect for $\\ell\\ge 3000.$ Weak lensing also mixes and distorts the EE and BB polarisation power spectra \\cite{zaldarriaga&seljak}. Finally, weak lensing also correlates with other secondary effects, such as the integrated Sachs-Wolfe (ISW) effect % and the Sunyaev-Zel'dovich (SZ) effect, % thus inducing non-Gaussianities which are manifest in higher-order point correlation functions \\cite{zaldarriaga,cooray}.  The relevance of the CMB lensing reconstruction for cosmology is threefold. First, by converting a fraction of the dominant E-mode polarisation due to density perturbations into a B-mode polarisation, CMB lensing introduces a contaminant in the measurement of the primordial B-mode polarisation and consequently of the inflation energy scale % \\cite{gravitational_waves}. Second, CMB lensing is also a powerful cosmological probe of the matter distribution integrated from the last scattering surface at redshift $z\\sim 1000$ to the present time. Third, by changing the power spectra of the perturbations and inducing non-Gaussianities, CMB lensing affects cosmological parameter estimation. Delensing the CMB will thus allow one to recover the primordial B-mode, probe the full-sky large-scale structure distribution with a maximum efficiency at $z\\approx 3$ \\cite{lewis&challinor}, and obtain unbiased parameter constraints from CMB measurements \\cite{lewis}.   Effort has been devoted to developing an optimal reconstruction of the lensing potential in harmonic space, which implicitly assumes full-sky coverage without galactic cuts, bad pixels due to excision of point sources or nonuniform weighting for uneven sky coverage \\cite{okamoto&hu}. In this idealized context, it has been shown how to construct an optimal quadratic estimator for a reconstruction based on the temperature anisotropy alone \\cite{hu} (see \\cite{hu&okamoto} for an extension of the formalism to include the polarisation). The improvement gained from using a maximum likelihood, and thus more optimal, estimator for the reconstruction from the temperature is marginal \\cite{hirata&seljak2002} (although this might not necessarily be the case for experiments sensitive to the B-mode polarisation \\cite{hirata&seljak2003}). This is because the distortion due to lensing is small compared to the intrinsic cosmic variance or to the experimental noise.   Our interest lies in considering a slightly less optimal estimator modified to have a finite range in real space. The estimator is based on a convolution of the square of the lensed temperature map with a local kernel, as has been suggested previously \\cite{bucher}. An estimator with kernel in real space acts locally, as opposed to the conventional estimator with kernel in harmonic space which acts over the entire celestial sphere even though it has finite support in harmonic space. The extent in real space will be the inverse of that in harmonic space, thus finite and determined by the scale of the root mean square of the lensing % convergence, with the contribution from sources beyond the kernel being negligible. An estimator with these properties can easily be adapted to include cuts, excisions of pixels and nonuniform sky coverage for a local extraction of the lensing % convergence. In particular, a nonuniform sky coverage can be accounted for by a corresponding weighting in the real space convolution.  The lensing distortion of the CMB anisotropy on the surface of last scattering can be described in three ways: by the lensing potential $\\psi,$ by the deflection vector $\\boldsymbol{\\alpha}=\\nabla \\psi$ or by the convergence tensor $\\boldsymbol{\\kappa}=-\\nabla\\nabla\\psi/2.$ The descriptions of $\\psi$ and $\\boldsymbol{\\alpha}$ suffer from an ambiguity upon translation, since a patch of the sky and its translation have the same likelihood on account of isotropy. In contrast, the description of the convergence, which is a gradient of the deflection vector field, is locally well defined. Here we implement the real space estimator to reconstruct the convergence % and compare the reconstruction to that obtained with the harmonic space estimator.  The convergence tensor can be decomposed as a sum of an isotropic (diagonal) tensor and an anisotropic (traceless) tensor as follows \\ba \\boldsymbol{\\kappa} =\\left( \\begin{array}{cc} \\kappa_{0}+\\kappa_{+} & \\kappa_{\\times}\\\\ \\kappa_{\\times} & \\kappa_{0}-\\kappa_{+} \\end{array} \\right). \\ea For a small patch of sky, we can use the flat-sky approximation and define a Cartesian coordinate system with basis vectors $\\{\\boldsymbol{e}_{x}(r),\\boldsymbol{e}_{y}(r)\\}$ and metric $g={\\rm diag}(1,1).$ In particular, the convergence $\\kappa_{0}=-(\\partial_{x}^2+\\partial_{y}^2)\\psi/2$ magnifies a feature on the last-scattering surface, the shear $\\kappa_{+}=-(\\partial_{x}^2-\\partial_{y}^2)\\psi/2$ stretches it along the $x$--axis while compressing it along the $y$--axis, and analogously the shear $\\kappa_{\\times}=-\\partial_{x}\\partial_{y}\\psi$ stretches it along the $y=x$ axis and compresses it along the $y=-x$ axis. Here we focus on the dilation effect captured by the isotropic component $\\kappa_{0},$ and will treat the shear components in a forthcoming study.  The convergence can be further decomposed into two contributions, namely the convergence of the unlensed temperature anisotropies, to a high degree isotropic, and the convergence of the lensing field, by nature random. In order to remove the zero mode of the unlensed CMB, we average the convergence  on each local patch over the sky region surveyed and subtract this average value from the estimated convergence map. Higher multipoles of the unlensed CMB will still be present in the estimated convergence map, which give rise to a noise bias. This bias can be removed by cross-correlating the reconstructed map with maps of the large scale structure derived from independent methods, e.g. a galaxy redshift survey. It can also be removed when estimating the convergence power spectrum by combining, in harmonic space, different triangle configurations with a common side. The formulation of an estimator for the lensing convergence given the lensed image still requires some prior knowledge of the statistics of the underlying temperature field, in particular the power spectrum of the unlensed sky which is assumed to be Gaussian.  This is the first of a series of studies where we explore the properties of the proposed estimator and test its implementation. The manuscript is organized as follows. In Section~\\ref{sec:weak_lensing} we % describe the weak lensing CMB of photons in the Born approximation. In Section~\\ref{sec:real_estimator} we derive the proposed estimator defined in real space, whose kernel in harmonic space is described in Appendix~\\ref{app:harmonic_estimator_description}. In Section~\\ref{sec:reconstruction} we proceed to reconstruct the convergence power spectrum by implementing the real space estimator on lensed maps synthesised as detailed in Appendix~\\ref{app:map_synthesis}. For comparison, we also implement the harmonic space estimator % as detailed in Appendix~\\ref{app:harmonic_estimator_implementation}. We discuss the results obtained from the two estimators in Section~\\ref{sec:discussion} and conclude  in Section~\\ref{sec:summary} by discussing further developments of the real space estimator in follow-up studies. ",
        "conclusions": "\\label{sec:discussion}  For each experiment we plot the output power spectrum $C_{\\ell}^{\\kappa_{out}\\kappa_{out}}$, both before (light grey solid line) and after (dark grey solid line) the removal of the bias, against the input power spectrum $C_{\\ell}^{\\kappa_{in}\\kappa_{in}}$ (black solid line). In Figs.~\\ref{fig:cl_kk2_designer} and \\ref{fig:cl_kk2_planck} we plot the reconstructed power spectra for the {\\designer} and the {\\planck} experiment respectively. We also plot the variance of the estimator ${\\rm Var}[\\hat\\kappa_{0}(\\boldsymbol{\\ell})]$ (black dot-dashed line) as well as % ${\\rm Var}[\\hat\\kappa_{0}(\\boldsymbol{\\ell})]/f_{sky}$ (grey dot-dashed line) as an auxiliary tool for visualization. The error bars on the recovered power spectra are computed from the standard deviation of the reconstructed power spectrum, binned over logarithmically spaced intervals in $\\ell$ space and added in quadrature. The variance of the reconstructed power spectrum encompasses both the variance of the optimal estimator and the average value of the convergence fluctuation in the map, weighted by the fraction of the area of the sky covered by the patch, $f_{sky},$ as follows \\cite{tegmark} \\ba {\\rm Var}[C_{\\ell}^{\\kappa_{0}\\kappa_{0}}] \\approx{1\\over {\\ell~f_{sky}}}\\left( {\\rm Var}[\\hat \\kappa_{0}(\\boldsymbol{\\ell})]+C_{l}^{\\kappa_{0}\\kappa_{0}}\\right)^2. \\ea We observe that the degree of recovery of the input power spectra is not the same throughout the interval in $\\ell.$ In particular, at small $\\ell$ the cosmic variance dominates the error of the reconstruction, whereas at large $\\ell$ the variance of the estimator takes over. The region where $C_{\\ell}^{\\kappa_{in}\\kappa_{in}}> {\\rm Var}[\\hat\\kappa_{0}(\\boldsymbol{\\ell})]$ indicates the interval in $\\ell$ where the lensing potential can in theory be mapped. % Accordingly, this is also the region where the reconstructed convergence power spectra tracks more closely the input power spectrum. Also visible are the fluctuations in the recovered power spectrum which arise from the subtraction of the noise bias term estimated from a finite number of realizations of the unlensed CMB map.  For the {\\planck} experiment, using the real space implementation we achieve a reasonable reconstruction of the input power spectrum in the interval where the input signal is larger than the variance of the estimator, although on smaller scales, $\\ell > 600,$ there is a decrease in the recovered power. Comparing the reconstructed power spectra for the {\\designer} and {\\planck} experiments we observe that the real space implementation is fairly insensitive to the detector noise and returns a reconstruction consistent with the input power spectrum. We can understand these results as follows.  \\begin{figure}[t] \\setlength{\\unitlength}{1cm} \\vskip-1.5cm \\centerline{ \\hskip-0.5cm \\includegraphics[width=12cm]{% pikassoyla5aHSdesigner.pdf} \\hskip-3cm \\includegraphics[width=12cm]{% pikassoyla5bRSdesigner.pdf} } \\vskip-7.5cm \\caption{\\baselineskip=0.5cm{ {\\bf Mean square of the convergence map reconstructed with the harmonic space estimator (left panel) and the real space estimator (right panel).} The smooth black curve is the input power spectrum. The light grey and dark grey solid lines are the output power spectra before and after the removal of the bias respectively. The error bars measure the total standard deviation binned over logarithmically spaced intervals in $\\ell.$ Here $\\theta_{fwhm}=7.8',$ $FOV_{map}=66.6^{\\circ}$ and $m_{max}=4,$ $\\ell_{max}=4000$ for no detector noise. }} \\label{fig:cl_kk2_designer} \\end{figure}  \\begin{figure}[h] \\setlength{\\unitlength}{1cm} \\vskip-1.5cm \\centerline{ \\hskip-0.5cm \\includegraphics[width=12cm] {% pikassoyla6aHSplanck.pdf} \\hskip-3cm \\includegraphics[width=12cm] {% pikassoyla6bRSplanck.pdf} } \\vskip-7.5cm \\caption{\\baselineskip=0.5cm{ {\\bf Mean square of the convergence map reconstructed with the harmonic space estimator (left panel) and the real space estimator (right panel).} The smooth black curve is the input power spectrum. The light grey and dark grey solid lines are the output power spectra before and after the removal of the bias respectively. The error bars measure the total standard deviation binned over logarithmically spaced intervals in $\\ell.$ Here $\\theta_{fwhm}=7.8',$ $FOV_{map}=66.6^{\\circ},$ and $m_{max}=4,$ $\\ell_{max}=4000$ for PLANCK noise at $\\nu=143~{\\rm GHz}.$ }} \\label{fig:cl_kk2_planck} \\end{figure}  As attested by the power spectra of the maps reconstructed with the real space estimator, there is a loss of power at small scales beginning around the angular scale corresponding to the size of the kernel. This loss of power is a consequence of averaging modes smaller than the finite extent of the kernel and hence is an intrinsic constraint of the real space implementation. This intrinsic constraint, studied in more detail in Ref.~\\cite{bucher2010}, can be overcome by shrinking the extent of the kernel in real space, which can be achieved with a smaller beam and detector noise, thereby moving the support for the kernel to larger $\\ell.$ The $\\ell$ interval over which the power spectrum can be recovered by the real space estimator is determined by the characteristic lengths of the map and the kernel as follows. The lower $\\ell$ limit is determined by the size of the map, since it measures the longest wavelength mode that can be enclosed within the size of the map. The upper $\\ell$ limit is determined by the size of the kernel, since it measures the smallest wavelength mode that the kernel can probe, below which there is a loss of power due to the averaging of modes smaller than the size of the kernel. Within this $\\ell$ range, the real space estimator seems to perform fairly well both in the absence and in the presence of detector noise. Note that the $\\ell$ range for a reasonable  reconstruction with the harmonic space estimator is wider, being limited at higher $\\ell$ by the angular scale corresponding to the beam size, since this is the scale which constrains the action of the kernel in harmonic space.  Comparing the reconstructed power spectra in the absence and in the presence of detector noise, we observe that the real space estimator appears to be insensitive to the experimental noise. This is because the estimate of the convergence in each pixel is given by the sum of the product of pairs of neighbouring pixels weighted by the kernel. As a result, the noise, being independent in each pixel, is averaged out. To substantiate this result, we reconstructed the convergence map from a pure white noise input map using the real space estimator and the harmonic space estimator. In Fig.~\\ref{fig:cl_kk2_planck_pure_noise} we observe that the power spectrum recovered by the real space estimator is about seven orders of magnitude smaller than the input power spectrum, whereas the power spectrum recovered by the harmonic space estimator is comparable to the input power spectrum.  Despite our attempt to remove the CMB and the detector noise bias from the reconstructed convergence power spectrum, there still remains a bias in the harmonic space implementation, as observed in Figs.~\\ref{fig:cl_kk2_designer} and \\ref{fig:cl_kk2_planck}. This bias presumably arises from a coupling between unlensed temperature modes due to the finite size of the map, providing an additional source of lensing which we have not treated here. This mode coupling is absent in the reconstruction over the full sky \\cite{inpaint}. The real space implementation appears to be insensitive to this bias.     \\begin{figure}[t] \\setlength{\\unitlength}{1cm} \\vskip-1.5cm \\centerline{ \\hskip-0.5cm \\includegraphics[width=12cm] {% pikassoyla7noise.pdf} } \\vskip-7.5cm \\caption{\\baselineskip=0.5cm{ {\\bf Mean square of the convergence map reconstructed with the harmonic space estimator (light grey line) and the real space estimator (dark grey line) from a pure white noise map.} The grey solid lines are the output power spectra whereas the smooth black curve is the input power spectrum. Here $\\theta_{fwhm}=7.8',$ $FOV_{map}=66.6^{\\circ},$ and $m_{max}=4,$ $\\ell_{max}=4000$ for PLANCK noise at $\\nu=143~{\\rm GHz}.$ }} \\label{fig:cl_kk2_planck_pure_noise} \\end{figure}    In this paper we implemented a new estimator of the convergence field for the extraction of the lensing potential from CMB maps. Our interest was in reconstructing the convergence field from CMB maps from which contaminants, such as point sources and the SZ effect, had previously been removed using multi-frequency information.  The new estimator acts locally in real space, thus being able to treat the excision of pixels and nonuniform sky coverage in a flexible manner. We implemented the estimator on two experimental setups, one without and one with detector noise, based on the specifications of  PLANCK for the $\\nu=143~{\\rm GHz}$ channel. For comparison of the performance of the new estimator, we also implemented the conventional estimator which acts in harmonic space. From the two experiments we learned that by increasing the range of the kernel in $\\ell$ space, and consequently reducing its range in $\\theta$ space, we improve the reconstruction of the power spectrum at small scales. The finite extent of the kernel is a desirable property of the proposed estimator, since in theory it allows the reconstruction of the lensing convergence, manifested essentially at very small scales, from a small map of the sky as long as the kernel probes sufficiently small angular scales. Even though there is a loss of power on scales smaller than the finite extent of the kernel, this effect could be studied using simulations to determine a form factor that could be applied to correct the loss of power \\cite{bucher2010}. We also observed that although the real space estimator is limited, in a statistical sense, to be as good as the harmonic space estimator, it has the advantage over the harmonic space estimator of being less sensitive to white noise. In practice, however, CMB experiments have to deal with correlated noise, which will pose as much of a challenge to the real space estimator as to the harmonic space estimator, most likely requiring the construction of a more carefully designed kernel in harmonic space that downweights the correlated modes before transforming to real space.   On the basis of the results presented here, we envisage two follow-up studies. In the first study we intend to optimize the current implementation of the real space estimator to use a kernel capable of probing smaller angular scales % in the reconstruction of the lensing convergence, as we would expect for the ACT and SPT experiments. As we have demonstrated, the real space estimator can be applied to small patches of sky without incurring in the serious spectral leakage that affects the harmonic space estimator on rectangular, rather than torodial, domains in the presence of red spectra. In real space, the local filter acts up to a kernel length away from the edge of the map without leakage of power. We also intend to study the effect of a mask due to the excision of bad pixels, e.g. arising from the removal of bright point sources, that can cause aliasing of power from large scales to small scales. Inpainting  has been proposed as a means of interpolating across masked regions that preserves the statistical properties of the map \\cite{inpaint}, so we plan to investigate how sensitive the real space estimator is to complex masks with and without the use of such techniques.  A straightforward extension will be to develop the analogous estimator for the shear components of the convergence tensor and work out how to combine the different components for the reconstruction of the lensing potential. Since the shear components do not contain the $\\ell=0$ mode, the corresponding estimators will be less sensitive to a bias. The dilatation and shear components provide complementary information on the lensing potential but are not independent, being related by consistency conditions that are algebraic in harmonic space. Since these components probe the monopole and quadrupole sectors respectively, the noise associated with these components will be uncorrelated and the reconstruction can be harmonized, with the noise being reduced via inverse variance weighting of these components. This point is discussed in more detail in Ref.~\\cite{bucher2010}.  In the second study we intend to extend the implementation of the real space estimator to CMB polarisation so as to optimize the reconstruction of the lensing potential from the PLANCK data.   \\vspace{0.6cm} \\centerline{\\bf"
    },
    "1005/1005.3527_arXiv.txt": {
        "abstract": "We report on the discovery and the timing analysis of the first eclipsing accretion-powered millisecond X-ray pulsar (AMXP): SWIFT~J1749.4--2807. The neutron star rotates at a frequency of $\\sim$517.9~Hz and is in a binary system with an orbital period of 8.8 hrs and a projected semi-major axis of $\\sim$1.90 lt-s. Assuming a neutron star between 0.8 and 2.2 $M_{\\odot}$ and using the mass function of the system and the eclipse half-angle, we constrain the mass of the companion and the inclination of the system to be in the $\\sim$0.46-0.81 M$_{\\odot}$ and $\\sim74.4^\\circ-77.3^\\circ$ range, respectively. To date, this is the tightest constraint on the orbital inclination of any AMXP. As in other AMXPs, the pulse profile shows harmonic content up to the 3rd overtone.  However, this is the first AMXP to show a 1st overtone with rms amplitudes between $\\sim6$\\% and $\\sim23$\\%, which is the strongest ever seen, and which can be more than two times stronger than the fundamental. The fact that SWIFT~J1749.4--2807 is an eclipsing system which shows uncommonly strong harmonic content suggests that it might be the best source to date to set constraints on neutron star properties including compactness and geometry. ",
        "introduction": "\\label{sec:intro}  The first accreting millisecond X-ray pulsar (hereafter AMXP) was discovered in 1998 \\citep[SAX~J1808.4--3658, see][]{Wijnands98} and since then, a total of 13 AMXPs have been found and studied in detail \\citep{Patruno10a}.  Most AMXPs show near sinusoidal profiles during most of their outbursts.  This is consistent with a picture in which only one of the hotspots (at the magnetic poles) is visible (see ref. below).  Deviations from a sinusoidal profile (i.e, an increase in harmonic content) are generally interpreted as being caused by the antipodal spot becoming visible, perhaps as accretion rate falls and the disk retreats \\citep[see, e.g.][ and references therein]{Poutanen03,Ibragimov09}.  Although the amplitude of the 1st overtone may reach that of the fundamental late in the outburst \\citep[see, e.g][]{Hartman08, Hartman09a}, no AMXP so far has shown pulse profiles where the 1st overtone is generally stronger than the fundamental throughout the outburst.   The stability of the pulse profiles in some of the AMXPs means that pulse profile modeling can be used to set bounds on the compactness of the neutron star and hence the dense matter equation of state \\citep[see, e.g.][ and references therein]{Poutanen03,Poutanen09}. Unfortunately, there is often a large degeneracy between the parameters due to the number of free parameters needed to construct the model profile. One of these parameters is the inclination of the system, which to date has not been well-constrained for any AMXP.    \\begin{figure}[http] \\center \\resizebox{1\\columnwidth}{!}{\\rotatebox{0}{\\includegraphics{./rmsvstime.eps}}} {\\footnotesize \\caption{\\textit{Top panel:} 2-10 keV flux as measured from RXTE/PCA and Swift/XRT observations. The flux of the last PCA observation (MJD 55307.5) is not shown; the spectrum of this observation was used as an estimate background emission (see text). Upper limits are quoted at 95\\% confidence level.  We calculated the flux during the last PCA and last Swift/XRT detection using WebPIMMS (assuming a power law spectrum with index 1.8). \\textit{Middle panels:} Fractional rms amplitude and 95\\% confidence level upper limits of the fundamental and three overtones as a function of time. Detections ($>3\\sigma$ single trial) and upper limits are from $\\sim500$ and $\\sim3000$ sec datasets, respectively. \\textit{Bottom panel:} Ratio between the fractional rms amplitude of the 1st overtone and fundamental. Blue triangles represent points in which both harmonics are significantly detected in $\\approx500$ second datasets.  Black circles represent the ratio between the fractional rms amplitude of the 1st overtone and the 95\\% confidence level upper limit to the amplitude of the fundamental. Grey circles represent the same ratio but when the fundamental is significantly detected and not the 1st overtone. This means that black circles represent lower limits while grey circles are upper limits. These ratios are independent of the background.}} \\label{fig:lc} \\end{figure}  In this Letter we report on the discovery and timing of the accretion-powered millisecond X-ray pulsar SWIFT~J1749.4--2807. Thanks to the observed eclipses \\citep{Markwardt10}, we set the tightest constraint on system inclination for any AMXP.  This, coupled with the fact that the amplitude of the first overtone is higher/comparable to that of the fundamental for much of the outburst and that the amplitude of the first overtone is unusually high, allows to put tight constraints on pulse profile models.  We show that SWIFT~J1749.4--2807 has the potential to be one of the best sources for this approach to constraining the neutron star mass-radius relation and hence the EoS of dense matter. ",
        "conclusions": ""
    },
    "1005/1005.0263_arXiv.txt": {
        "abstract": "It is shown that the Hubble constant can be derived from the standard luminosity function of galaxies as well as from a new luminosity function as deduced from the mass-luminosity relationship for galaxies. An analytical expression for the Hubble constant can be found from the maximum number of galaxies (in a given solid angle and flux) as a function of the redshift. A second analytical definition of the Hubble constant can be found from the redshift averaged over a given solid angle and flux. The analysis of two luminosity functions for galaxies brings to four  the new definitions of the Hubble  constant. The equation that regulates the Malmquist bias for galaxies is derived and as a consequence it is possible to extract a complete sample. The application of these  new formulae to the data of the two-degree Field Galaxy Redshift Survey provides a Hubble constant of $( 65.26 \\pm 8.22 ) \\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1}$ for a redshift lower than 0.042. All the results  are deduced in a Euclidean universe because the concept of space-time curvature is not necessary as well as in a static universe because two mechanisms for the redshift of galaxies alternative to the Doppler effect are invoked. ",
        "introduction": "The Hubble constant, in the following $H_0$, is defined as \\begin{equation} H_0 = \\frac{v}{D} [\\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1}] \\quad , \\end{equation} where $v=cz$ is the recession velocity, $D$ is the distance in $Mpc$, $c$ is the velocity of light and $z$ is the redshift defined as \\begin{equation} z = \\frac { \\lambda_{obs} - \\lambda_{em} } { \\lambda_{em}} \\quad , \\end{equation} with $\\lambda_{obs}$  and $\\lambda_{em}$ denoting respectively the wavelengths of the observed and emitted lines as  determined from the lab source. The first numerical values of the Hubble constant were : $H_{0}=625$ $ \\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1}$ as deduced by \\citet{Lemaitre1927}, $H_{0}=460$ $ \\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1}$ as deduced by \\citet{Robertson1928}, $H_{0}=500 $ $\\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1}$ as deduced by \\citet{Hubble1929}  and $H_{0}=290$ $ \\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1}$ as deduced by \\citet{Oort1931}. Figure~\\ref{hystory} reports the decrease of the numerical value of the Hubble constant from 1927 to 1980.  \\begin{figure*} \\begin{center} \\includegraphics[width=10cm]{f01.eps} \\end {center} \\caption { Logarithmic values of the Hubble constant $H_0$ from 1927 to 1980. The error bar is evaluated according to the file http://www.cfa.harvard.edu/~huchra/hubble.plot.dat~. } \\label{hystory}% \\end{figure*} At the time of writing, two excellent reviews have been written, see \\citet{Tammann2006} $( H_{0}=(63.2\\pm1.3~ (random) ~\\pm5.3~ (systematic)) $ $ \\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1} ) $ and \\citet{Jackson2007} ($ H_0\\sim 70\\mbox{\\,--\\,}73$ $ \\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1} $). We now report the methods that use the global properties of galaxies as indicators of distance: \\begin{enumerate} \\item Luminosity classes of spiral galaxies; $H_{0}=(55 \\pm 3) $ $\\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1}$ , see \\citet{Sandage1999b} \\item 21 cm line widths; $H_{0}=(59.1 \\pm 2.5)$ $ \\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1}$ , see \\citet{Federspiel1999} \\item Brightest cluster galaxies; $H_{0}=(54.2 \\pm 5.4 )$ $ \\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1}$ , see \\citet{Sandage1973} \\item The D$_{n}$-$\\sigma$ or fundamental plane method; $H_{0}=(57 \\pm 4)$ $ \\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1}$ , see \\citet{Federspiel1999}  \\item Surface brightness fluctuations; $H_{0}=71.8 $ $\\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1}$ , see \\citet{Tammann2006} \\item Gravitational lens; $H_{0}=(72\\pm 12) $ $\\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1}$ , see \\citet{Saha2006} \\item The Sunyaev--Zel'dovich effect; $H_{0}=(67\\pm 18)$ $ \\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1}$ , see \\citet{Udomprasert2004} \\item Ks-band Tully-Fisher Relation; $H_{0}=(84 \\pm 6 )$ $ \\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1}$ , see \\citet{Russell2009}, where the Hubble constant was named Hubble parameter.  \\end {enumerate}  At the time of writing, the first important evaluation of the Hubble constant is through Cepheids (key programs with HST) and type Ia Supernovae, see \\citet{Sandage2006}, \\begin{equation} H_0 =(62.3 \\pm 5 ) \\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1} \\quad . \\label{h0cefeidi} \\end {equation} A second  important evaluation comes from the  three years of  observations with the Wilkinson Microwave Anisotropy Probe, see Table 2 in \\citet{Spergel2007}; \\begin{equation} H_{0}=(73.2 \\pm 3.2) \\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1} \\quad . \\label{hzerowmap} \\end{equation} In the  following,  we will process galaxies  having redshifts as given by the catalog of galaxies. The forthcoming analysis is based on two key assumptions: (i) the flux of radiation from galaxies in a given wavelength decreases with the square of the distance; (ii) the redshift is assumed to have a linear relationship with distance in $Mpc$. These two hypotheses allow some new physical mechanisms to be accepted which produce a linear relationship between redshift and distance, for redshifts lower than 1. In this framework, we can speak of a Euclidean universe because the distances are deduced from the  Pythagorean theorem and a static universe because it is not expanding. The already listed approaches leave a series of questions unanswered or partially answered: \\begin {itemize} \\item Can the Hubble constant be deduced from the Schechter luminosity function of galaxies? \\item Can the Hubble constant be deduced from a new luminosity of galaxies alternative to the Schechter function? \\item Can the equation that regulates the Malmquist bias be derived in order to deal with a complete sample in apparent magnitude? \\item Can the reference magnitude of the sun be deduced from the luminosity function of galaxies? \\end{itemize} In order to answer these questions, Section~\\ref{secprelimaries} contains three introductory paragraphs on sample moments, the weighted mean and the determination of the so-called \"exact value\" of the Hubble constant. Section~\\ref{useful} reviews the basic system of magnitudes, a review of two alternative mechanisms for the redshift of galaxies, two analytical definitions of the Hubble constant in terms  of the Schechter luminosity function of galaxies and two other definitions that can be found by adopting a new luminosity function for galaxies. Section~\\ref{secnumerical} contains a numerical evaluation of the four new formulae for the Hubble constant as deduced from the data of the two-degree Field Galaxy Redshift Survey. Section \\ref{secmsun} contains a numerical evaluation of the reference magnitude of the sun for a given catalog. ",
        "conclusions": "A careful study of the standard LF of galaxies allows the determination of the position of the maximum in the theoretical number of galaxies versus redshift and the theoretical averaged redshift. From the two previous analytical results, it is possible to extract two new formulae for the Hubble constant, equations~(\\ref{hzero1}) and (\\ref{hzero2}). The same procedure can be applied by analogy to a new LF as given by the mass-luminosity relationship, see equations~(\\ref{hzero3}) and (\\ref{hzero4}). The weighted mean of the four  values of $H_0$ as deduced from Table~\\ref{hubblevalue} gives  \\begin{equation} H_0 =( 65.26 \\pm 8.22  ) \\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1} \\quad when~ z \\leq 0.042 \\nonumber \\quad . \\end {equation} This value lies between the value deduced from the Cepheids, see~\\citet{Sandage2006} and formula~(\\ref{h0cefeidi}) and the value deduced from WMAP,  see~\\citet{Spergel2007} and formula~(\\ref{hzerowmap}).  The developed framework also enables the deduction of the reference magnitude of the sun, see formula~(\\ref{msun}) and the application to the 2dFGRS gives \\begin{equation} M_{\\sun} = (5.5 \\pm 0.35) \\quad . \\end{equation}  Assuming that the exact value is $ M_{\\sun}$ = 5.33, the efficiency in deriving the reference magnitude of the sun is $\\epsilon=96.63 ~\\% $ when $H_0=66.04 \\mathrm{\\ km\\ s}^{-1}\\mathrm{\\ Mpc}^{-1}$. We briefly review the basic cosmological assumptions adopted here to derive the Hubble constant: \\begin{itemize} \\item The mechanism that produces the redshift, here extracted from the catalog of galaxies, is not specified but we remember that the plasma redshift and DET (Dispersive Estinction Theory) do not produce a geocentric model for the universe as given by the Doppler shift, see \\citet{Wang2007}.  \\item The number of galaxies as a function of redshift as well as the averaged redshift are evaluated in a Euclidean space or, in other words, the effects of space-curvature are ignored. \\item The spatial inhomogeneities present in the catalog of galaxies are partially neutralized by  the operation of adding together the data of the south and north galactic pole of the 2dFGRS. The transition from a nonhomogeneous to a quasi-homogeneous universe is clear when Figure~\\ref{isto1} and Figure~\\ref{isto2} are carefully analyzed. \\item The  initial  assumptions of: (i) natural  flux decreasing as given by equation~(\\ref{flr2}) ; (ii) linear relationship between redshift and distance which are present in the joint distribution in {\\it z} and {\\it f} for the number of galaxies  are  justified by the acceptable results obtained for the theoretical maximum in the number of galaxies, see Figure~\\ref{zeta_max_flux}. This fact allow us to speak of a Euclidean universe up to $z \\leq 0.042$. \\item The presence of the Malmquist bias does not allow to extrapolate the concept of a Euclidean, static universe for distances greater than $z\\,>\\,0.042$ when the 2dFGRS  catalog is considered. \\end{itemize}"
    },
    "1005/1005.4489_arXiv.txt": {
        "abstract": "By observing mergers of compact objects, future gravity wave experiments would measure the luminosity distance to a large number of sources to a high precision but not their redshifts. Given the directional sensitivity of an experiment, a fraction of such sources (gold plated -- GP) can be identified optically as single objects in the direction of the source. We show that if an approximate distance-redshift relation is known then it is possible to statistically resolve those sources that have multiple galaxies in the beam. We study the feasibility of using gold plated sources to iteratively resolve the unresolved sources, obtain the self-calibrated best possible distance-redshift relation and provide an analytical expression for the accuracy achievable. We derive lower limit on the total number of sources that is needed to achieve this accuracy through self-calibration.  We show that this limit depends exponentially on the beam width and give estimates for various experimental parameters representative of future gravitational wave experiments DECIGO and BBO. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.3056_arXiv.txt": {
        "abstract": "{Passive early-type galaxies (ETGs) provide an ideal laboratory for studying the interplay between dust formation around evolved stars and its subsequent destruction in a hot gas. Using {\\it Spitzer}-IRS and {\\it Herschel} data we compare the dust production rate in the envelopes of evolved AGB stars with a constraint on the total dust mass. Early-type galaxies which appear to be truly passively evolving are not detected by {\\it Herschel}. We thus derive a distance independent upper limit to the dust grain survival time in the hostile environment of ETGs of $< 46\\pm 25\\;\\rm Myr$ for amorphous silicate grains. This implies that ETGs which \\emph{are} detected at far-infrared wavelengths have acquired a cool dusty medium via interaction. Given likely time-scales for ram-pressure stripping, this also implies that only galaxies with dust in a cool (atomic) medium can release dust into the intra-cluster medium. } ",
        "introduction": "In this paper we compare the dust outflow rate from the evolved AGB stars in passive ETGs, with the total dust mass, to estimate dust grain lifetimes. Both quantities are derived from observations; the outflow rate from {\\it Spitzer}-IRS spectra, and the total dust mass from {\\it Herschel} (Pilbratt et al. 2010) maps of the {\\it Herschel} Virgo Cluster Survey, HeViCS (Davies et al. 2010, www.hevics.org). We assume that evolved stars are the only source of dust in these passive systems as there is very little evidence that type Ia supernovae produce dust.   IRAS data have been used in the past in a similar way. Soifer et al. (1986) estimated the mass-loss from evolved stars in the bulge of M~31 from 12 and $25\\;\\rm \\mu m$ fluxes and the total dust mass from the 60 and $100\\;\\rm \\mu m$ fluxes. They concluded that given that the observed dust mass could be produced by stellar outflows in only $10^7\\;\\rm yr$, there must be some mechanism that depletes the bulge of inter-stellar matter. Jura et al. (1987) made similar calculations for a sample of elliptical galaxies. They estimated dust outflow rates by assuming that half of the $12\\;\\rm \\mu m$ flux came from dusty stellar envelopes. They also found that the stellar mass-loss rate would be sufficient to produce the observed cool ISM in much less than a Hubble time.  Dust mass production rate and total dust mass can only reasonably be compared in objects that are truly passively evolving. Temi et al. (2007) find, that for a given blue luminosity, the 70 and $160\\;\\rm \\mu m$ luminosities vary by 2 orders of magnitude, indicating that many ETGs have a significant dust component not directly attributable to mass-loss from evolved stars.  For our purposes, these ETGs are not `passive' because other processes (mergers?)  have probably contributed to the dust mass.  Recent studies of ETGs with {\\it Spitzer} (Bressan et al. 2006; Clemens et al. 2009) have shown that even in samples of ETGs selected to be the most passive objects (ie. lying on the colour-magnitude relation) a significant fraction show evidence of either on-going or recent past star formation. These results indicate that ETGs in which dust features are seen in the optical (e.g. Sadler \\& Gerhard, 1985) and many that have been detected by IRAS at 60 and $100\\;\\rm \\mu m$ (Knapp et al. 1989) and {\\it Spitzer} at 70 and $160\\;\\rm \\mu m$ (Kaneda et al. 2007) may also host low levels of star formation. Intriguigly, of the 7 ellipticals detected at 70 and $160\\;\\rm \\mu m$ by Kaneda et al. (2007) the only object not showing PAH emission features is actually a radio source, so that the far-infrared emission is probably synchrotron!  Tsai and Mathews (1995) studied dust destruction in ETGs via thermal sputtering in the hot ($10^6 - 10^7$\\,K),  low density ($n_H \\sim 0.1$\\,cm$^{-3}$) gas,  and found that the destruction time-scale is short compared to any cooling  flow or dust transport time-scale. They therefore concluded that dust is destroyed `on the spot' before it has time to migrate within a galaxy. In a following work, Tsai and Mathews (1996) found dust-to-gas mass ratios that are orders of magnitude less than that in typical stellar ejecta or in the ISM of the Milky Way. Their interpretation of the IRAS data was, at that time, hampered by the resolution of the existing observations; they were therefore not able to tie down the origin and distribution of the observed dust emission.  As part of the {\\it Herschel} Science Demonstration Phase (SDP), a $4\\times4\\;\\rm deg^2$ field, centred approximately on M~87, has been observed as part of the HeViCS with both the PACS (Poglitsch et al. 2010) and SPIRE (Griffin et al. 2010) instruments. See Davies et al. (2010) for details. ",
        "conclusions": "We interpret the ratio of dust mass to dust production rate as a measure of grain lifetime. In principle, however, dust could be removed from the galaxy rather than destroyed.   A wind of hot gas produced by type Ia supernovae probably cannot drive dust (or gas) out of the potential well of massive ETGs (Mathews \\& Loewenstein 1986). However, good evidence of ram-pressure stripping of the hot ISM in elliptical galaxies has been found in the form of X-ray tails (Randall et al. 2008; Machacek et al. 2006). If dust is found in this hot gas then it too may be removed. However, the timescale for such stripping is estimated to be of the order a few $10^8\\;\\rm yr$ (Takeda et al. 1984; Murakami \\& Babel 1999, eq. 12) which is an order of magnitude longer than our upper limit to the grain lifetime. The observation that the X-ray emission of most ETGs is \\emph{not} displaced from the stellar distribution, in fact, suggestes that either the stripping time-scale is longer than the hot gas production time-scale or that the ram-pressure is typically insufficient to remove the hot medium. It is now clear that if dust is isolated from the hot medium in denser, cooler clouds, than ram-pressure does have time to strip dust (Cortese et al. 2010) but the passive ETGs in our study do not show any evidence of accumulating dust in such clouds. \\emph{Passive} ETGs do not directly pollute the intra-cluster medium with dust.   Grain lifetimes of $<46\\;\\rm Myr$ have immediate implications for ETGs that \\emph{are} detected in the FIR. Either they produce dust at rates much higher than those given in Table~\\ref{table:1} or their dust is shielded from the hot gas in cool clouds. For example, the dust mass of NGC~4435 of $1.2 \\times 10^{6}\\;\\rm M_{\\odot}$ would require a total dust production rate of $0.03\\;\\rm M_{\\odot}\\,yr^{-1}$, more than an order of magnitude greater than any value in Table~\\ref{table:1}. This object is known to host star formation (Panuzzo et al., 2007) and so dust is likely to be in cool clouds more typical of late-type galaxies; grain lifetimes in this case may be much longer. Dust may have been acquired from an interaction, but in any case it probably arrived as part of a cooler medium.   In a passive ETG, the ISM, and the dust within it, come mainly from mass-loss from evolved stars. Unless the dust-to-gas ratio in this released material is very variable, one would expect dust and gas production to follow one another. If the X-ray luminosity, $L_{x}\\propto n_{e}^2$ one would naively expect a relation $L_{x}\\propto M_{\\rm d}^2$. However, if the grain lifetime, $t_{\\rm g}\\propto 1/n_{\\rm H}$, then we would expect only a linear correlation between $L_{x}$ and $M_{\\rm d}$. Kaneda et al. (2007) find evidence of an \\emph{anti-correlation} for a small sample of elliptical galaxies detected by {\\it Spitzer}-MIPS, but this may be expected if the dust were of external origin.   Only NGC~4371 is detected by MIPS at $70\\;\\rm \\mu m$ (marginally at $160\\;\\rm \\mu m$). The detection is consistent with our {\\it Herschel} detection limits. This object actually shows extremely weak PAH features in its {\\it Spitzer}-IRS spectrum, and is therefore probably not totally passive. The object is included for comparison, and its IRS spectrum included in Fig.~\\ref{fig:1}.  We have assumed that AGB stars are the only source of dust. Were type Ia supernovae able to produce dust from the metals they produce, we estimate that the contribution to the dust production rate would be, at most, similar to that of AGB stars. Our limit on grain lifetime would be a factor of $\\sim$2 lower in this case.   Although dust grains are rather short-lived in ETGs, the prospects for survival in the intra-cluster  medium are rather better because the gas densities are typically at least 2 orders of magnitude lower than in ETGs (B{\\\"o}hringer et al., 1994). Therefore, if dust can be removed from cluster galaxies, or form in the intra-cluster medium, there is a good chance that such dust will be detected by the HeViCS, especially at larger cluster radii where gas densities are lower."
    },
    "1005/1005.1924_arXiv.txt": {
        "abstract": "Intergalactic magnetic fields (IGMF) can cause the appearance of halos around the gamma-ray images of distant objects because an electromagnetic cascade initiated by a high-energy gamma-ray interaction with the photon background is broadened by magnetic deflections. We report evidence of such gamma-ray halos in the stacked images of the 170 brightest active galactic nuclei (AGN) in the 11-month source catalog of the \\textit{Fermi} Gamma-Ray Space Telescope. Excess over point spread function in the surface brightness profile is statistically significant at $3.5\\sigma$ (99.95\\% confidence level), for the nearby, hard population of AGN. The halo size and brightness are consistent with IGMF, $B_{\\rm IGMF}\\approx 10^{-15} ~ {\\rm G}$. The knowledge of IGMF will facilitate the future gamma-ray and charged-particle astronomy. Furthermore, since IGMF are likely to originate from the primordial seed fields created shortly after the Big Bang, this potentially opens a new window on the origin of cosmological magnetic fields, inflation, and the phase transitions in the early Universe. ",
        "introduction": "Intergalactic magnetic fields (IGMF) had not been measured until now, despite their importance for gamma-ray and cosmic-ray astronomy and their likely connection to the primordial fields that could have seeded the stronger magnetic fields observed in galaxies, Sun, and Earth. This is because IGMF are too small for conventional astronomical probes, such as Zeeman splitting or Faraday rotation. Unlike the fields in galaxies, which are believed to have been amplified by the dynamo action of the large-scale convective motions of gas, the fields in voids remain low, close to their primordial values modified only by the relatively small contribution of the fields leaking out of galaxies \\citep{Kronberg1994,Grasso2001,Widrow2002,Kulsrud2007}. The observational and theoretical upper bounds on IGMF constrain their magnitudes to be below $10^{-9}$~G \\citep*{1997PhRvL..78.3610B}, whereas any value above $\\sim$10$^{-30}$~G is sufficient to explain the $\\sim \\mu$G Galactic magnetic fields generation by the dynamo mechanism \\citep*{1999PhRvD..60b1301D}.   One can detect such extremely weak fields using high-energy gamma rays \\citep*{Aharonian1994,Plaga1995}. Very energetic photons emitted from active galactic nuclei (AGN) or other strong sources produce pairs of electrons and positrons in their interactions with the extragalactic background light (EBL). These pairs up-scatter the cosmic microwave background photons to high energies, giving rise to an electromagnetic cascade, and the photons from the cascade are detected by gamma-ray telescopes, such as {\\it Fermi}. Since the trajectories of electrons and positrons in the cascade are affected by magnetic fields, a gamma-ray image of AGN is expected to exhibit a halo of secondary photons around a bright central point-like source \\citep{Aharonian1994,Dolag2009,Neronov2009}. The central image is expected to be composed of photons emitted directly from the source with energies below the pair production threshold. In addition, delays in arrival times of the secondary photons can be used to probe IGMF \\citep{Plaga1995,Ando2004,Murase2008}. Finally, at TeV energies, the secondary photons produced in interactions of cosmic rays with EBL may have already been observed by the air Cherenkov telescopes \\citep{Essey2010a,Essey2010b}.   Thus far, in TeV range, HEGRA \\citep{HEGRA} and MAGIC \\citep{MAGIC} did not detect any halo component of two bright blazars, Mrk~501 and Mrk~421, and they set upper limits on the flux. In particular, the analysis of MAGIC using gamma rays above 300~GeV excludes some range of IGMF between $4\\times 10^{-15}$ and $10^{-14}$~G. Very recently, IGMF above $3\\times 10^{-16}$~G were proposed as an explanation of non-observation by {\\it Fermi} of several AGN known to be bright TeV sources (\\citealt{Neronov2010}; see also \\citealt{2010MNRAS.tmpL..82T}).  In this {\\it Letter}, we present evidence of extended images and of IGMF at $3.5\\sigma$ level, based on gamma-ray data collected by the Large Area Telescope (LAT) onboard {\\it Fermi}, in the energy range between 1~GeV and 100~GeV. It is consistent with pair-halo scenario with IGMF, $B_{\\rm IGMF} \\approx 10^{-15} ~ {\\rm G}$. The knowledge on IGMF will facilitate the future gamma-ray and cosmic ray astronomy, and it will open a new window on the origin of cosmological magnetic fields \\citep{Cornwall:1997ms}, inflation \\citep{Turner:1987bw,DiazGil:2007dy},  and the phase transitions in the early Universe \\citep{Vachaspati:1991nm,Baym:1995fk,Vachaspati:2001nb}. ",
        "conclusions": ""
    },
    "1005/1005.3921_arXiv.txt": {
        "abstract": "We use a suite of cosmological, hydrodynamical simulations to investigate the chemical enrichment history of the Universe. Specifically, we trace the origin of the metals back in time to investigate when various gas phases were enriched and by what halo masses. We find that the age of the metals decreases strongly with the density of the gas in which they end up. At least half of the metals that reside in the diffuse intergalactic medium (IGM) at redshift zero (two) were ejected from galaxies above redshift two (three). The mass of the haloes that last contained the metals increases rapidly with the gas density. More than half of the mass in intergalactic metals was ejected by haloes with total masses less than $10^{11}\\,\\Msun$ and stellar masses less than $10^9\\,\\Msun$. The range of halo masses that contributes to the enrichment is wider for the hotter part of the IGM. By combining the `when' and `by what' aspects of the enrichment history, we show that metals residing in lower density gas were typically ejected earlier and by lower mass haloes. ",
        "introduction": "Despite the fact that `metals' - heavy elements produced by stars - make up a very tiny portion of the cosmic matter budget, they are of critical importance to our understanding of the Universe. From a diagnostic point of view, they are very useful in constraining star formation and tracing feedback from massive stars and supernovae. More importantly, however, metals impact the rate at which gas cools, affecting structure formation on scales of galaxy clusters down to dust grains. This effect should not go understated since, depending on the temperature, metals can reduce the cooling time by over an order of magnitude for solar metallicity \\cite[e.g.][]{Sutherland1993, Gnat2007, Wiersma2009a}.  Observations clearly show that a substantial fraction of the diffuse intergalactic medium (IGM) has been enriched with heavy elements \\citep[e.g.][]{Cowie1998,Schaye2000a,Ellison2000,Schaye2003,Simcoe2004,Aracil2004,Schaye2007,Aguirre2008,Danforth2008,Cooksey2010,Pieri2010}. However, the physical mechanism, the timing and the sources of the enrichment all remain unknown. The IGM may have been enriched by the first generations of stars and galaxies at very high redshifts, or it may have been polluted by more massive galaxies at intermediate redshifts. Metals may be carried out into intergalactic space by galactic winds driven by supernovae, radiation pressure from star light, or by active galactic nuclei (AGN) \\citep[e.g.][]{Cen1999a,Aguirre2001,Aguirre2001z3,Madau2001,Theuns2002,Tornatore2004,Aguirre2005,Oppenheimer2006,Dave2007,Oppenheimer2008,Scannapieco2006b,Kawata2007,Kobayashi2007,Sommer-Larsen2008,Wiersma2009b,Shen2009,Tornatore2010,Choi2010}. In dense environments such as groups and clusters of galaxies metals may also be efficiently dispersed via processes such as ram pressure stripping or tidal stripping (see e.g.\\ \\citealt{Borgani2008} for a review).   The most straightforward way to investigate when the IGM received its metals, is to plot the metallicity or metal mass as a function of redshift, as we did for a series of cosmological simulations in \\citet{Wiersma2009b}. By comparing models with different physical implementations,  one can obtain information about the processes that are important for IGM enrichment. However, the metals that reside in the IGM at high redshift may not be the same metals that are in the IGM at the present day. What is hot IGM gas at high redshift, may have cooled by low redshift. Intergalactic metals may also accrete onto galaxies and end up in the ISM or in stars.  Here we also use cosmological simulations to study the enrichment of the cosmos - noting that the IGM is of particular interest - but we take a novel approach in that we classify the gas based on its state at the redshift of interest (we choose $z=0$ and $z=2$) and trace its enrichment history back in time with the aim of understanding how the observed metals were put into place. Specifically, we will address the following two questions: When did gas that resides in the IGM at $z=0$ and $z=2$ receive its metals? And what are the masses of the haloes that last contained these metals?  This paper is organised as follows. In Section~\\ref{sec-method} we describe the simulations used, emphasising the details that are particularly pertinent here. Section~\\ref{sec-when} discusses \\textit{when} the gas was enriched, taking each sub-phase in turn. In Section~\\ref{sec-what} we investigate \\textit{what} enriched the gas, using a halo-finder to determine what kinds of objects are mainly responsible for the enrichment of various gas phases. We connect the two ideas in Section~\\ref{sec-conn}, relating enrichment time with the last halo that contained the metals that end up in the IGM. Finally, we summarise our findings in Section~\\ref{sec-summary}. ",
        "conclusions": "\\label{sec-summary}  We have used cosmological, hydrodynamical simulations that include radiative cooling, star formation, stellar evolution and supernova feedback, and, in some cases, AGN feedback, to investigate the metal enrichment history of the Universe. Specifically, we took advantage of the Lagrangian nature of our simulations to investigate when gas that ends up in a particular gas phase at redshift $z=0$ or $z=2$ received its metals and what the masses were of the haloes that last contained these metals.  We considered different physical models, all taken from the OWLS project, including simulations with a different implementation of supernova feedback and a model that includes AGN feedback. While there are some small differences, such as slightly higher enrichment redshifts for models with more efficient feedback in massive galaxies, the main trends are strikingly similar for all models. We cannot rule out that some of our conclusions may depend on our prescriptions for feedback or that they are specific to SPH. However, the fact that three drastically different feedback models give such similar results suggests that our conclusions are robust.  The time since the enrichment varies most strongly with the gas density. Metals in lower density gas are typically much older than metals in high-density gas, a trend that extends over ten orders of magnitude in gas density. At least half of the metals that reside in the diffuse IGM at $z=0$ ($z=2$) were ejected above redshift two (three).  The enrichment redshift also varies with the temperature of the gas, but this mostly reflects the fact that the typical gas density of enriched gas varies with its temperature. For gas that is dense enough to be able to cool, the time since enrichment correlates well with the cooling time. Gas with shorter cooling times received its metals more recently.  The typical mass of the haloes from which the metals residing in the IGM were ejected, increases rapidly with the gas density. At least half of the metal mass was ejected by haloes with total masses less than $10^{11}\\,\\Msun$, which corresponds to stellar masses smaller than $10^9\\,\\Msun$. For the low-redshift IGM the mass of the dominant IGM polluters may be substantially smaller than this because our predictions for present day low-density gas are not close to converged with respect to the numerical resolution.  The age of the metals ending up in the diffuse IGM is strongly anti-correlated with the mass of the haloes from which they were ejected. In other words, older metals were typically ejected by lower mass haloes. This anti-correlation is less strong for the hotter part of the IGM (the WHIM), for which recent enrichment by low-mass haloes is more important. While massive haloes ($> 10^{12}\\,\\Msun$) are unimportant for the enrichment of the diffuse IGM, they do contribute to the pollution of the warm-hot IGM at low redshift.  Our main and strongest conclusion is that metals residing in lower density gas were typically ejected earlier and by lower mass haloes. This suggests that travel-time is a limiting factor for the enrichment of the IGM, as proposed by \\citet{Aguirre2001z3}. Metals that have been ejected recently simply have not had sufficient time to reach the low-density gas far from galaxies. The anti-correlation between metal age and halo mass would then follow because typical galaxy masses are lower at higher redshifts. Another mechanism that could explain the trends is fall back onto galaxies. If metals ejected by lower mass galaxies are less likely to fall back \\cite[e.g.][]{Oppenheimer2010} or if they typically leave the haloes with larger velocities, then we would expect lower mass galaxies to be more important for gas that is further away. Finally, low-mass galaxies are more weakly clustered than high-mass galaxies, thus low-density gas is more accessible to low-mass galaxies.  It is not immediately obvious how one could compare this result with observations. One way would be to consider abundance ratios of elements in various phases as a function of density and redshift in an attempt to find a signature that could be searched for in high-redshift observations. However, uncertainties in the nucleosynthetic yields and the type Ia supernova rates \\citep[see][]{Wiersma2009b}, as well as in the ionization corrections, may make such a comparison troublesome. Correlations between absorbers and the distances to different types of galaxies may also help shed light on the enrichment mechanism \\citep[e.g.][]{Steidel2010,Chen2009,Wilman2007,Stocke2006}, although is is important to keep in mind that the metals that are observed to be near a galaxy may have been injected by its lower-mass progenitors \\citep{Porciani2005,Scannapieco2005}. In fact, our results would imply that this is likely to be the case for metals in low-density gas. We leave an analysis of abundance ratios and other observational comparisons to a future work.  The importance of low-mass galaxies is a challenge for simulations of the enrichment of the IGM, as it implies the need for high-resolution in large volumes. On the other hand, our results suggest that the low-density IGM provides us with an exciting opportunity to study the consequences of feedback in low-mass galaxies, a key but poorly understood ingredient of models of galaxy formation, and that it provides us with a fossil record of galaxy formation in the high-redshift Universe."
    },
    "1005/1005.0958_arXiv.txt": {
        "abstract": "Quasi-equilibrium sequences of binary neutron stars are constructed for a variety of equations of state in general relativity. Einstein's constraint equations in the Isenberg-Wilson-Mathews approximation are solved together with the relativistic equations of hydrostationary equilibrium under the assumption of irrotational flow. We focus on unequal-mass sequences as well as equal-mass sequences, and compare those results. We investigate the behavior of the binding energy and total angular momentum along a quasi-equilibrium sequence, the endpoint of sequences, and the orbital angular velocity as a function of time, changing the mass ratio, the total mass of the binary system, and the equation of state of a neutron star. It is found that the orbital angular velocity at the mass-shedding limit can be determined by an empirical formula derived from an analytic estimation. We also provide tables for 160 sequences which will be useful as a guideline of numerical simulations for the inspiral and merger performed in the near future. ",
        "introduction": "Coalescing binary neutron stars are among the most promising sources of gravitational waves for ground-based laser-interferometric gravitational-wave detectors such as LIGO \\citep{bro04}, GEO600 \\citep{luc06}, TAMA300 \\citep{and05}, and VIRGO \\citep{ace07}. Merger of binary neutron stars, together with that of black hole-neutron star binaries, is also considered to be one of the candidates for the central engines of short-hard gamma-ray bursts \\citep{nar92}. These facts motivate us to study coalescing binary neutron stars.  Binary neutron stars evolve as a result of gravitational radiation reaction and eventually merge. This evolutionary sequence is usually divided into three stages, depending on the characteristic timescales associated with orbital period and gravitational radiation reaction, as well as on the tidal effects for each neutron star. The first stage is the adiabatic, inspiral phase. In this phase, the timescale of orbital shrink due to the emission of gravitational waves is much longer than the orbital period, and thus, the binary system evolves adiabatically. In addition, each neutron star can be treated as a point mass, because the orbital separation is much larger than the neutron star radius and hence the tidal deformation of the neutron star is negligible.  In this phase, a post-Newtonian approximation together with the point particle approximation is a robust tool for determining the orbital evolution and for computing gravitational waveforms (see e.g. \\citet{bla06} and references therein).  The second stage is called the intermediate phase or the quasi-equilibrium phase. In this phase, the binary system is considered to be still in the adiabatic, inspiral phase, but we need to take into account tidal effects on each neutron star, i.e., hydrodynamic effects in neutron stars, as well as full effects of general relativity, because the orbital separation between two neutron stars is only a few times of the neutron-star radius and thus they are in a strong two-body gravitational field. One of the aims of the present paper is to contribute to the understanding of this phase. We will explain more details about the purpose of the present paper later.  The last stage is the merger phase, for which the timescale of orbital shrink becomes shorter than the orbital period and thus the evolution of the system proceeds in a dynamical manner. Furthermore, the system becomes highly general relativistic, because the compactness of the system, defined by the ratio of the gravitational radius to the radius of the system, becomes larger than $\\sim 0.2$. To clarify the merger phase, numerical relativity is the unique approach. Since the first fully general relativistic merger simulation was performed by \\citet{shi99}, huge effort has been devoted in this research field \\citep{shi00,shi02,shi03,shi05,shi06a,due03,mil04,and08a,and08b,yam08,liu08,bai08,gia09,kiu09,kiu10}. (See e.g. \\citet{oec07a} and \\citet{oec07b} for works focusing on the micro-physics in a neutron star but not in fully general relativistic framework.)  Now we return to the intermediate phase and explain the purposes for studying the quasi-equilibrium phase of binary neutron stars in general relativity in detail. There are two roles for this study. One is to clarify the physical conditions in this phase, for example, how large the tidal deformation of a neutron star is, when the mass-shedding from the neutron star occurs, and what the orbital angular velocity at the mass-shedding limit is. The other is to provide initial data for studying the merger phase by numerical relativity simulations. Numerical relativity, in which Einstein's evolution equations are solved, requires initial data that satisfy Einstein's constraint equations and also that should be as physical as possible. Obviously, it is important to derive accurate initial data for a scientific study. Constructing such initial data is just obtaining relativistic binary neutron stars in quasi-equilibrium.  The first effort on this issue was devoted to constructing corotating binary neutron stars in general relativity because implementing a numerical code for computing such solutions is relatively easy \\citep{bau97,bau98,mar98,usu00,usu02,tan02b,tan03,tic09}. \\citet{koc92} and \\citet{bil92}, however, found that the timescale of coalescence driven by the gravitational-radiation reaction is much shorter than that of synchronization due to the viscosity in a neutron star. This implies that if the spin angular velocity of neutron stars is much smaller than the orbital angular velocity in a late inspiral phase, we can regard the rotation state of a neutron star to be approximately irrotational for the subsequent phase until the merger sets in. Additionally, any neutron star spins down due to electromagnetic radiation during its life from birth to the coalescence. The spin down timescale of a neutron star in a known binary is at longest as short as the coalescing timescale \\citep[larger than $\\sim 10^8$ yrs;][]{lor08}. Moreover, the spin period of neutron stars in a known binary is always larger than 20 ms which is $\\sim 10$ times larger than the orbital period in the late inspiral phase just prior to merger, 2--3 ms. Therefore, we can conclude that the irrotational flow is physically more realistic.\\footnote[2]{If a first born neutron star in a binary was strongly recycled during the evolution of the companion, it may result in fast rotation and weak magnetic field. In such a case, the neutron star may rotate on the order of milliseconds even just before the merger. Some effort to approximately construct quasi-equilibrium, non-corotating, and non-irrotational binary systems is reported in \\citet{mar03} and \\citet{bau09}.}  Under the assumption of irrotation, formulation for solving relativistic hydrostatic equations was derived \\citep{bon97,asa98,shi98,teu98}. Soon after the formulation was derived, quasi-equilibrium sequences of irrotational binary neutron stars were calculated \\citep{bon99,mar99,ury00a,ury00b,gou01,tan02b,tan03,bej05}, based on the Isenberg-Wilson-Mathews (IWM) approximation to general relativity \\citep{ise78,ise08,wil89}. (See \\citet{shi04} for an advanced formulation and \\citet{ury06,ury09} for the results.) Even though a lot of sequences have been, so far, calculated, systematic survey has not yet been done. In particular, {\\it unequal-mass}, irrotational binary neutron stars with an {\\it equation of state other than single polytrope} has not been studied in detail. In \\citet{tan02b,tan03}, quasi-equilibrium sequences of unequal-mass binaries were calculated, but a polytropic equation of state (EOS) was used. In \\citet{bej05} and \\citet{ury09}, non-polytropic EOSs were used, but quasi-equilibrium sequences of non-equal-mass binaries were not computed. Actually, unequal-mass, irrotational binary neutron stars with realistic EOSs in quasi-circular orbits were constructed and used initial data for merger simulations in \\citet{shi05}, \\citet{shi06a}, and \\citet{kiu09,kiu10}.\\footnote[3]{We use the term ``realistic equations of state'' for the EOSs derived from nuclear physics, although no one really knows a realistic one.} We, however, have not constructed sequences, and rather computed only some initial data sets for each neutron star mass. The purpose of the present paper is to complete the issue and to provide a database of the sequences.  To compute unequal-mass binary systems of arbitrary mass ratio, we develop a new code for the present research, because the numerical code that we developed for the previous works \\citep{tan02b,tan03} had a problem with calculating binary systems composed of significantly different-mass neutron stars in general relativity even though the problem was not in Newtonian computation \\citep{tan02a}. As we will explain in Section 2, the method to determine the center of mass of unequal-mass binary systems, i.e., the position of the rotation axis, caused the problem in the previous code, but we have overcome it by employing a new method.  In addition, we employ a wide variety of EOSs; piecewise polytropic EOSs \\citep{rea09a,rea09b}, tabulated realistic EOSs derived from nuclear physics, and fitted EOSs to the tabulated realistic EOSs. Some of the first and second EOSs were, respectively, used in \\citet{ury09} and in \\citet{bej05}, but we adopt a wider set of EOSs in this paper. Furthermore, we systematically study the unequal-mass binaries, whereas \\citet{bej05} and \\citet{ury09} focused only on the equal-mass case.  This paper is organized as follows. We briefly review the basic equations and explain the improvement of the numerical code in Section 2. In Section 3, the results for the code test are shown. In Section 4, we show numerical results and discuss the effects of EOS on each sequence. Section 5 is devoted to a summary. Throughout this paper we adopt geometrized units with $G=c=1$, where $G$ denotes the gravitational constant and $c$ the speed of light. Latin and Greek indices denote purely spatial and spacetime components, respectively. ",
        "conclusions": "In the present paper, quasi-equilibrium sequences of binary neutron stars are constructed for 18 EOSs except for the $\\Gamma=2$ polytrope, in the IWM framework in general relativity. The EOSs we choose are nine piecewise polytropes, six tabulated realistic EOSs derived by using various theories of dense nuclear matter and different solution methods of the many-body problem in nuclear physics, and three EOSs expressed by a fitting formula. We employ three total masses, $M_0=2.4 M_{\\odot}$, $2.7 M_{\\odot}$, and $3.0 M_{\\odot}$ for each EOS, and compute sequences of three mass ratios for each total mass: i.e., $M_{\\rm ADM}^{\\rm NS1}$ versus $M_{\\rm ADM}^{\\rm NS2}$ $=1.20 M_{\\odot}$ versus $1.20 M_{\\odot}$, $1.10 M_{\\odot}$ versus $1.30 M_{\\odot}$, and $1.00 M_{\\odot}$ versus $1.40 M_{\\odot}$ for $M_0=2.4 M_{\\odot}$; $1.35 M_{\\odot}$ versus $1.35 M_{\\odot}$, $1.25 M_{\\odot}$ versus $1.45 M_{\\odot}$, and $1.15 M_{\\odot}$ versus $1.55 M_{\\odot}$ for $M_0=2.7 M_{\\odot}$; and $1.50 M_{\\odot}$ versus $1.50 M_{\\odot}$, $1.40 M_{\\odot}$ versus $1.60 M_{\\odot}$, and $1.30 M_{\\odot}$ versus $1.70 M_{\\odot}$ for $M_0=3.0 M_{\\odot}$.  We focus on the unequal-mass sequences and compare their results with those of the equal-mass case. Changing the mass ratio, the total mass, and the EOSs, we investigate the behavior of the binding energy and total angular momentum along a sequence, the endpoint of sequences, and the orbital angular velocity as a function of time. For example, it is found for the piecewise polytropic EOSs that the change in stellar radius fixing the stiffness of the core EOS makes the orbital angular velocity at the mass-shedding limit vary widely, while the change in stiffness of the core EOS fixing the stellar radius does not change the orbital angular velocity at the mass-shedding limit significantly.  It is also found that the orbital angular velocity at the closest separation decreases as we decrease the mass ratio, $M_{\\rm ADM}^{\\rm NS1}/M_{\\rm ADM}^{\\rm NS2} \\le 1$. The reason is that the less massive star in an unequal-mass binary is tidally deformed by the companion more massive star and starts shedding mass at larger separation than that for the equal-mass case. It is found that the orbital angular velocity at the mass-shedding limit increases as the neutron-star mass increases. This is because a more massive star becomes more compact and more difficult to be tidally disrupted for the same EOS. This implies that the binary neutron stars with massive stars need to come closer than those with less massive stars for reaching the mass-shedding limit. The orbital angular velocity at the mass-shedding limit is analyzed by using a Newtonian argument, and an empirical formula is found as \\begin{equation} M_0 \\Omega_{\\rm ms} =0.270~ {\\cal C}_{\\rm NS1}^{3/2} \\Bigl( 1+\\frac{1}{q} \\Bigr)^{3/2} q^{1/2}. \\end{equation}  We have provided tables for 160 sequences as shown in Appendix \\ref{app:sequence}. Those tables may be useful as one of the database for future works on binary neutron stars in quasi-equilibrium and as a guideline of numerical simulations for the inspiral and merger."
    },
    "1005/1005.1390_arXiv.txt": {
        "abstract": "We present a new theoretical calculation of the contribution to the extragalactic gamma-ray background radiation (EGRB) from star-forming galaxies, based on a state-of-the-art model of hierarchical galaxy formation that is in quantitative agreement with a variety of observations of local and high-redshift galaxies. Gamma-ray luminosity ($L_\\gamma$) and spectrum of galaxies are related to star formation rate ($\\psi$), gas mass ($M_{\\rm gas}$), and star formation mode (quiescent or starburst) of model galaxies using latest observed data of nearby galaxies.  We try the two limiting cases about gamma-ray production: the escape limit ($L_\\gamma \\propto \\psi M_{\\rm gas}$) and the calorimetric limit ($L_\\gamma \\propto \\psi$), and our standard model predicts 7 and 4\\% contribution from star-forming galaxies to the total EGRB flux (including bright resolved sources) recently reported by the {\\it Fermi} Gamma-Ray Space Telescope. Systematic uncertainties do not allow us to determine the EGRB flux better than by a factor of $\\sim$ 2. The predicted number of nearby galaxies detectable by {\\it Fermi} is consistent with the observation. Intergalactic absorption by pair-production attenuates the EGRB flux only by a modest factor of $\\sim$1.3 at the highest {\\it Fermi} energy band, and the reprocessed cascade emission does not significantly alter EGRB at lower photon energies. The sum of the known contributions from AGNs and star-forming galaxies can explain a large part of EGRB, with a remarkable agreement between the predicted model spectrum and observation. ",
        "introduction": "\\label{sec:INTRODUCTION} The existence of the extragalactic diffuse gamma-ray background (EGRB) has been revealed first by the SAS-2 satellite (\\citealt{1977ApJ...217L...9F}; \\citealt*{1978ApJ...222..833F}). Better determinations of the flux and spectrum of EGRB were achieved by the EGRET detector on board the Compton Gamma-Ray Observatory (\\citealt{1998ApJ...494..523S}; \\citealt*{2004ApJ...613..956S}). The most reliable measurement of EGRB has very recently been reported based on the data of the {\\it Fermi} Gamma-Ray Space Telescope (\\citealt{fermi}), and the EGRB spectrum is well described by a single power-law with a photon index of $2.41 \\pm 0.05$ and the photon flux of about $(1.03 \\pm 0.17) \\times 10^{-5} \\;{\\rm photons\\;cm^{-2}\\;s^{-1}\\;sr^{-1}}$ above 100 MeV (\\citealt{PhysRevLett.104.101101}).  The origin of EGRB has been discussed for a long time and various sources have been discussed as possible contributors to EGRB, such as active galactic nuclei (AGNs, especially blazars), galaxy clusters, intergalactic shocks produced by structure formation, and dark matter annihilation (see, e.g., \\citealt{2007AIPC..921..122D} for a review).  Almost all of the known extragalactic gamma-ray sources are blazars, and their contribution to the EGRB has been intensively studied (e.g., Inoue \\& Totani 2009; Venters 2010; Abdo et al. 2010d, and references therein). However, star-forming galaxies should also be gamma-ray emitters by cosmic-ray interactions with interstellar medium (ISM) and interstellar radiation field (ISRF) (\\citealt{2000ApJ...537..763S}; \\citealt{2004ApJ...613..962S}), and there must be non-zero contribution to EGRB. This is obvious because we know that the Galactic disk is a strong source of diffuse gamma-rays, and gamma-rays from Large Magellanic Cloud (LMC) have been detected by EGRET (\\citealt{1992ApJ...400L..67S}).  Furthermore, gamma-ray emission in GeV--TeV from Small Magellanic Cloud (SMC, \\citealt{2010arXiv1008.2127T}) and two nearby starburst galaxies, M82 and NGC 253, have recently been discovered by H.E.S.S. (\\citealt{2009Sci...326.1080A}), VERITAS (\\citealt{2009Natur.462..770V}), and {\\it Fermi} (\\citealt{2010ApJ...709L.152A}). The purpose of this paper is to present a new estimate of the contribution from star-forming galaxies to EGRB.   There are a number of previous studies on this issue (Strong et al. 1976; Lichti et al. 1978; Dar \\& Shaviv 1995; Pavlidou \\& Fields 2002; Thompson et al. 2007; Bhattacharya \\& Sreekumar 2009; Ando \\& Pavlidou 2009; Lacki et al. 2010; Fields et al. 2010), and the estimates of the contribution to EGRB ranges $\\sim$10--50\\%\\footnote{ Here, the contribution is against the total (physical) extragalactic background photon flux ($>$100 MeV) including bright resolved sources. It is often calculated against the unresolved component of EGRB, but it depends on the flux sensitivity of a particular detector.  Throughout the paper, ``the contribution to EGRB'' is defined against the total EGRB flux, which we estimated from the resolved and unresolved components of the {\\it Fermi} data (Abdo et al. 2010d).  We used a photon index of 2.4 for resolved blazars (Abdo et al. 2010d) to extrapolate the resolved component from 200 to 100 MeV.}.  Most of these studies calculated gamma-ray luminosity ($L_\\gamma$) from star formation rate (SFR) of galaxies, because the cosmic-ray energy input is expected to be proportional to SFR. However, if escape of cosmic-rays from galaxies is significant, SFR cannot be used as a reliable indicator of $L_\\gamma$ (see \\S \\ref{sec:ModelingLg} in more details).  Several studies used infrared luminosity of galaxies as SFR indicators, but IR luminosity (especially at far-IR in early-type galaxies) is not a perfect SFR estimator because a significant amount of dust can be heated by ISRF from relatively old stars (e.g., Salim et al. 2009 and references therein).  Amount of interstellar gas should also be important to determine $L_\\gamma$, because the degree of cosmic-ray escape is affected by the target amount.  In fact, recent observations indicate that gamma-ray luminosity is nicely correlated with the product of SFR and gas mass in galaxies(Abdo et al. 2010a).  Pavlidou \\& Fields (2002) included gas mass in the theoretical prediction of EGRB using the data of the cosmic star formation history (CSFH). Galaxies in the universe is considered as a closed box containing stellar and gas mass in present-day galaxies, and CSFH is used to solve the time evolution of relative fractions of stars and gas. However this approach likely overestimates the gas mass contributing to gamma-ray production at high redshifts, because most of the present-day stellar mass is in the form of gas and assumed to contribute to gamma-ray emission.  In reality, only the gas in collapsed dark halos can contribute to cosmic-ray interactions in galaxies, which is a small fraction at high redshifts according to the structure formation theory. Recently, Fields et al. (2010) incorporated gas mass in galaxies using the Schmidt-Kennicutt relation in their calculation of EGRB.  In this case one must know galaxy size to relate SFR and $M_{\\rm gas}$, and a single value of galaxy size at each redshift was assumed as a simple model.  Here we present a new calculation of the contribution from star-forming galaxies in EGRB using a state-of-the-art theoretical model of galaxy formation in the framework of hierarchical galaxy formation, which is in quantitative agreement with a variety of observations at high redshifts as well as the local universe. An advantage of our approach compared with previous studies is that we can calculate both SFR and gas mass of individual galaxies at various redshits.  Furthermore, we utilize the information of the gamma-ray spectra recently observed for nearby starbursts, in addition to the standard spectrum of the Galactic diffuse emission, to predict the EGRB spectrum based on the galaxy formation model including both quiescent and starburst galaxy populations.  The paper is organized as follows. In the next section we describe our model calculations. We present the results on EGRB and statistics about the number of nearby galaxies detectable by {\\it Fermi} in \\S \\ref{sec:RESULTS}. After discussion on the uncertainties in our calculation (\\S \\ref{sec:Uncertainties}) and implications for the origin of EGRB (\\S \\ref{sec:EGRB-origin}), conclusions will be presented in \\S \\ref{sec:Conclusions}.  In this work, cosmological parameters of $\\Omega_{0}=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_{0}=70~{\\rm Mpc^{-1}~km~s^{-1}}$ are adopted. ",
        "conclusions": "\\label{sec:Conclusions}  We have presented a new calculation of EGRB from cosmic-ray interactions in star-forming galaxies, based on a state-of-the-art galaxy formation model in the framework of hierarchical structure formation. The galaxy formation model is quantitatively consistent with various observations at high redshifts as well as the local universe. Gamma-ray luminosities of galaxies are calculated by star formation rate and gas mass in model galaxies, with the relation $L_\\gamma \\propto (\\psi M_{\\rm gas})^{0.86}$ (the escape limit) or  $L_\\gamma \\propto (\\psi)^{1.2}$ (the calorimetric limit) which are calibrated by the recent observational data of nearby galaxies.  The predicted number of nearby galaxies that are detectable by {\\it Fermi} is consistent with the actual number observed so far.  We found that star-forming galaxies make 4--7\\% contribution to the total EGRB flux reported by {\\it Fermi} in our standard model. Combined with the contribution from blazars as estimated by the {\\it Fermi} data, more than $\\sim$50\\% of EGRB can be accounted for. If the soft power-law tail of CXB is extending from MeV to GeV region as expected from the MeV background data and theoretical considerations of AGN accretion disks (Inoue et al. 2008), additional $\\sim$20 \\% can be explained.  The combined spectrum by AGNs and star-forming galaxies is remarkably similar to the observed EGRB spectrum.  It should be noted that there is no free parameters that can be tuned to fit the obeserved EGRB spectrum in the present model; the blazar EGRB spectrum of IT09 is determined by the spectral templates of the blazar SED (spectral energy distribution) sequence, and that of the star-forming galaxy component in the present work by templates constructed by observed nearby galaxies. Therefore we conclude that a large part ($\\gtrsim$70\\%) of EGRB can be explained by reasonable sources of AGNs and star-forming galaxies.  Further examination is required to see whether the apparent $\\lesssim$30\\% residual of EGRB is mainly a result of modeling uncertainties, experimental/observational uncertainties in deriving the EGRB data, or significant contributions from completely different source populations.  Given the good spectral agreement of the {\\it Fermi} data and AGNs/star-forming galaxies, the rest of EGRB is also likely dominated by astrophysical sources accelerating particles, even if a completely diffrent population is responsible for it.  It is in contrast that a large part of the blazar component of EGRB will be resolved into discrete sources by the ultimate {\\it Fermi} sensitivity in the near future (IT09), while almost all of the star-forming galaxy component will remain unresolved because of the faintness of individual sources. Therefore, any exotic contribution like dark matter annihilation cannot be probed directly under the level of the star-forming galaxies, i.e., $\\sim$ 10\\% of the total EGRB flux.  Another approach such as utilizing anisotropy would be required to search for the signal under that level (see, e.g. Ando et al. 2007; \\citealt{2009PhRvL.102x1301S}, \\citealt{2009arXiv0912.1854H}, \\citealt{2010arXiv1005.0843C})   We assumed a simple empirical relation between gamma-ray luminosity, SFR, and gas mass of galaxies.  Only two spectral templates were used in the calculation. A next step would be to construct a more physical model of gamma-ray luminosity and spectrum based on a larger number of physical quantities (e.g., size and stellar radiation field in addition to SFR and gas mass), taking into account propagation of cosmic-rays and production processes of gamma-rays."
    },
    "1005/1005.3489_arXiv.txt": {
        "abstract": "Iron, the Universe's most abundant refractory element, is highly depleted in both circumstellar and interstellar environments, meaning it exists in solid form. The nature of this solid is unknown. In this Letter, we provide evidence that metallic iron grains are present around oxygen-rich AGB stars, where it is observationally manifest as a featureless mid-infrared excess. This identification is made using \\emph{Spitzer Space Telescope} observations of evolved globular cluster stars, where iron dust production appears ubiquitous and in some cases can be modelled as the only observed dust product. In this context, FeO is examined as the likely carrier for the 20-$\\mu$m feature observed in some of these stars. Metallic iron appears to be an important part of the dust condensation sequence at low metallicity, and subsequently plays an influential r\\^{o}le in the interstellar medium. We explore the stellar metallicities and luminosities at which iron formation is observed, and how the presence of iron affects the outflow and its chemistry. The conditions under which iron can provide sufficient opacity to drive a wind remain unclear. ",
        "introduction": "Iron is the sixth most abundant element in the Universe, and the most abundant refractory element. It is observed to be highly depleted in both interstellar and circumstellar environments \\citep{SB79,MvWLE05,DIRMV09}, and must therefore predominantly exist in an unknown solid form. Iron may be incorporated in other dust grains, primarily silicates, though these are usually iron-poor \\citep{GS99,KdKW+02}. Alternatively, iron may form a metallic condensate \\citep{KdKW+02,VvdZH+09}.   Dust grains form around evolved stars by condensation from the gas phase, either directly, or onto molecular `seeds'. Individual dust species can usually be identified by distinct infrared emission bands. In oxygen-rich environments, amorphous or crystalline forms of silicates, spinels and corundum are observed to form, whilst carbon-rich environments give rise to amorphous carbon, graphite and SiC. Metallic iron grains, however, produce featureless infrared emission which can be difficult to differentiate from other sources, particularly amorphous carbon dust. Iron has hitherto only been inferred as a likely component of oxygen-rich dusty winds \\citep{KdKW+02,VvdZH+09}, but never positively identified.  Previous observational studies \\citep{MvLD+09,BMvL+09} of giant stars in globular clusters have found a featureless contribution in addition to the flux emanating from the star's photosphere, which has been attributed to an unidentified circumstellar dust species. We herein show that this dust species is metallic iron. ",
        "conclusions": "Using mid-infrared spectra and photometry, we have shown that there is a considerable amount of unattributed infrared flux in a large selection of globular cluster giants. Through radiative transfer modelling, we deduce that this flux is most likely due to metallic iron grains forming in a truncated stellar wind, corroborated by a potential identification of FeO with the 20-$\\mu$m emission feature. The production of metallic iron seems to become more dominant at lower metallicity, suggesting that the dust condensation sequence is fundamentally different in metal-poor objects. Large-scale production of iron grains and iron oxide in AGB stars can explain iron depletion in the gas and solid phases of the post-AGB stars and planetary nebulae; as well as isotopic ratios in FeO grains in meteorites. While iron increases opacity in oxygen-rich winds, it remains unclear whether it can dominate driving of metal-poor winds. Future observations to determine outflow velocities in these stars should help determine the driver of these stellar winds.   \\vspace{5 mm}  This paper uses observations made using the \\emph{Spitzer Space Telescope}, operated by JPL, California Institute of Technology under NASA contract 1407 and supported by NASA through JPL (contract number 1257184). We thank Martha Boyer for her helpful comments."
    },
    "1005/1005.4867_arXiv.txt": {
        "abstract": "{}  {Aims: To provide a significantly improved probability distribution for the H-test for periodicity in X-ray and $\\gamma$-ray arrival times, which is already extensively used by the $\\gamma$-ray pulsar community. Also, to obtain an analytical probability distribution for stacked test statistics in the case of a search for pulsed emission from an ensemble of pulsars where the significance per pulsar is relatively low, making individual detections insignificant on their own. This information is timely given the recent rapid discovery of new pulsars with the Fermi-LAT t $\\gamma$-ray telescope. } {Methods: Approximately $10^{14}$ realisations of the H-statistic ($H$) for random (white) noise is calculated from a random number generator for which the repitition cycle is $\\gg 10^{14}$. From these numbers the probability distribution $P(>H)$ is calculated. } {Results: The distribution of $H$ is is found to be exponential with parameter $\\lambda=0.4$ so that the cumulative probability distribution $P(>H)=\\exp{(-\\lambda H)}$. If we stack independent values for $H$, the sum of $K$ such values would follow the Erlang-K distribution with parameter $\\lambda$ for which the cumulative probability distribution is also a simple analytical expression. }",
        "introduction": "When searching for a periodic signal in X-ray or $\\gamma$-ray arrival times dominated by noise, we may either perform a blind search for $\\gamma$-ray pulsars as demonstrated by the Fermi-LAT Collaboration (Abdo et al. \\cite{fermi1}), or, search for such a signal where the frequency parameters have been prescribed by contemporary radio data (Weltevrede \\cite{fermi2}). Following the folding of say $N$ arrival times $t_1, ...,t_N$ modulo the pulsar spin parameters, we arrive at a set of phases $\\theta_i$, $i=1, ...N$. However, in the case of blind searches, Atwood et al. \\cite{atwood} introduced a time differencing technique in which case the number of trial periods is significantly reduced.  de Jager, Swanepoel \\& Raubenheimer (\\cite{dsr}, hereafter DSR) reviewed the general class of {\\it Beran} statistics (Beran \\cite{beran}), from which the most general test statistics such as Pearson's $\\chi^2$, Rayleigh and $Z^2_m$ statistics are derived, and from within this class they derived the well known {\\it H-test} for X-ray and $\\gamma$-ray Astronomy.  The probability distribution of the {\\it H-test} statistic as given by DSR was derived from Monte Carlo simulations employing $\\sim 10^8$ simulations. The computational power and random number simulators on typical IBM machines during the 1980's had limited ranges of applicability and the {\\it H-test} suffered accordingly. For values of $H<23$ we found that the probability distribution was exponential with parameter $\\lambda=0.398$ (or 0.4), whereas a hard tail developed for $H>23$, which resulted in a significant compromise in sensitivity.  The old version of the {\\it H-test} probability distribution is already extensively used by e.g. the Fermi-LAT Collaboration for pulsar searches (e.g. Abdo et al. \\cite{fermi1} and Weltevrede et al. \\cite{fermi2}), and from this paper it will become clear that the significances assigned to pulsar detections (or non-detections) may be too conservative, so that some pulsars may be missed, especially when many trial periods are involved, so that large values of the $H$-statistic are required for a significant detection . In this Letter we notify the community that all previous published significances from the {\\it H-test} should be reassessed, based on the new improved distribution presented below. Before we do so, we first briefly review the origin and properties of the {\\it H-test}. ",
        "conclusions": "For $M=1$ we retrieve the well-known {\\it Rayleigh test}, with the exception that the parameter for the exponential distribution has been reduced from $\\lambda=0.5$ (for the {\\it Rayleigh test}) to $\\lambda=0.4$ for the $H-test$. This slight loss of sensitivity is the effect of the number of trials taken implicitly into account as a result of the search through $m$ within the $H-test$.  Finally, it is clear that the corrected distribution of the {\\it H}-statistic follows a simple exponential with parameter $\\lambda=0.4$ and evaluation of results for $H>23$ (i.e. the $10^{-4}$ significance level) would yield more significant results compared to the old distribution presented by DSR. For example, for $H=50$ a probability of $4\\times 10^{-8}$ is typically quoted in the literature, whereas the true probability for uniformity is actually $2\\times 10^{-9}$ - already a factor 20 smaller.   A Fermi-LAT example of the Vela pulsar ($>500$ MeV) shows clearly values for $M$ up to 20 for ``skymap'' significances $X=p\\sqrt{N} > 20$, whereas $H$ scales with $H\\propto X^2$ as expected for Beran-type tests. The scaling $H=1.9X^2$ can be used to predict $H$-test statistics for Vela-like pulsars above 500 MeV if we assume the excess on the skymap is all pulsed.  The hard tail of the distribution beyond $H>23$ presented by DSR probably arose from the repitition cycle of the random number generator used in those days, so that the same fluctuations at large $H$ values were repeated given the finite cycle length of the generator used. In this case we however used a generator with a cycle time much longer than $10^{14}$, in which case we did not see the repitition of outliers as a result of a finite cycle length. Confirmation of the possible break (i.e. downturn) in the probability distribution at $H>70$ requires extensive simulations beyond $4\\times 10^{14}$ and is beyond the scope of this paper.  \\begin{figure} \\centering \\includegraphics[width=8 cm]{h-statistic_2.EPS} \\caption{The distribution of the {\\it H-statistic} derived from $4\\times 10^{14}$ Monte Carlo simulations with the best fit model for the cumulative probability $P(>H)=\\exp{(-0.4H)}$ (solid line) and the old version of the probability distribution given by DSR shown as a dashed line.} \\label{fig2} \\end{figure}"
    },
    "1005/1005.1445_arXiv.txt": {
        "abstract": "\\vspace{1cm} \\centerline{\\bf ABSTRACT}\\vspace{2mm} Very recently, Verlinde considered a theory in which space is emergent through a holographic scenario, and proposed that gravity can be explained as an entropic force caused by changes in the information associated with the positions of material bodies. Then, motivated by the Debye model in thermodynamics which is very successful in very low temperatures, Gao modified the entropic force scenario. The modified entropic force (MEF) model is in fact a modified gravity model, and the universe can be accelerated without dark energy. In the present work, we consider the cosmological constraints on the MEF model, and successfully constrain the model parameters to a narrow range. We also discuss many other issues of the MEF model. In particular, we clearly reveal the implicit root to accelerate the universe in the MEF model. ",
        "introduction": "\\label{sec1}  Very recently, Verlinde~\\cite{r1} considered a theory in which space is emergent through a holographic scenario, and proposed that gravity can be explained as an entropic force caused by changes in the information associated with the positions of material bodies. In this scenario, Verlinde has successfully derived the Newton's law of gravitation, the Einstein equations, and the law of inertia, from the entropic point of view. In fact, the entropic force scenario is similar to the old idea of Jacobson~\\cite{r2}, but also beyond it in some sense. Similar entropic insight into gravity has also been made by Padmanabhan~\\cite{r3} independently and simultaneously.  Here we briefly mention some key points of the entropic force scenario following~\\cite{r1}. Motivated by Bekenstein's argument~\\cite{r4}, Verlinde postulated that the change in entropy near the holographic screen is linear in the displacement $\\Delta x$, namely, \\be{eq1} \\Delta S=2\\pi k_B\\frac{mc}{\\hbar}\\Delta x\\,, \\ee where $m$ is the mass of test particle, whereas $k_B$, $c$ and $\\hbar$ are Boltzmann constant, speed of light and the reduced Planck constant, respectively. The effective entropic force acting on the test particle due to the change in entropy obeys the first law of thermodynamics \\be{eq2} F\\Delta x=T\\Delta S\\,, \\ee where $T$ is the temperature. If one takes the Unruh temperature $T$ experienced by an observer in an accelerated frame whose acceleration is $a$, i.e., \\be{eq3} k_B T=\\frac{1}{2\\pi}\\frac{\\hbar a}{c}\\,, \\ee to be the temperature associated with the bits on the holographic screen, from Eqs.~(\\ref{eq1})---(\\ref{eq3}), it is easy to recover the second law of Newton \\be{eq4} F=ma\\,. \\ee Considering a sphere as the holographic screen, Verlinde assumed that the number of used bits on the holographic screen $N$ is proportional to the area $A=4\\pi r^2$, i.e., \\be{eq5} N=\\frac{Ac^3}{G\\hbar}\\,. \\ee According to the equipartition law of energy, the total energy inside the screen is \\be{eq6} E=\\frac{1}{2}N k_B T\\,. \\ee Of course, one can identifies $E$ with the mass $M$ inside the screen through \\be{eq7} E=Mc^2. \\ee From Eqs.~(\\ref{eq1}), (\\ref{eq2}), and (\\ref{eq5})---(\\ref{eq7}), one can recover the Newton's law of gravitation \\be{eq8} F=G\\frac{Mm}{r^2}\\,, \\ee where $G$ can be identified with the Newton constant now. From Eqs.~(\\ref{eq3}), (\\ref{eq4}) and (\\ref{eq8}), it is easy to find the gravitational acceleration \\be{eq9} g=\\frac{GM}{r^2}\\,, \\ee and the temperature \\be{eq10} T=\\frac{\\hbar}{k_B c}\\frac{g}{2\\pi}\\,. \\ee As shown in~\\cite{r1}, a relativistic generalization of the presented arguments directly leads to the Einstein equations. We strongly refer to the original paper~\\cite{r1} for great details.  Soon after Verlinde's proposal of entropic force, many relevant works appeared. For examples, Cai, Cao and Ohta~\\cite{r5}, Shu and Gong~\\cite{r6} derived the Friedmann equations from entropic force simultaneously. Smolin~\\cite{r7} derived the Newtonian gravity in loop quantum gravity. Li and Wang~\\cite{r8} showed that the holographic dark energy can arise in the entropic force scenario. Easson, Frampton and Smoot~\\cite{r9} considered the entropic accelerating universe and the entropic inflation. Tian and Wu~\\cite{r10}, Myung~\\cite{r11} discussed the thermodynamics of black holes in the entropic force scenario. Vancea and Santos~\\cite{r12} considered the uncertainty principle from the point of view of entropic force. Zhang, Gong and Zhu~\\cite{r13}, Sheykhi~\\cite{r14} derived the modified Friedmann equation from the corrected entropy. Also, Modesto and Randono~\\cite{r15} discussed the corrections to Newton's law from the corrected entropy. Cai, Liu and Li~\\cite{r16} considered a unified model of inflation and late-time acceleration in the entropic force scenario. For other relevant works to entropic force, we refer to e.g.~\\cite{r17,r18,r19,r36} and references therein.  The works mentioned above are in fact closely following Verlinde's proposal of entropic force~\\cite{r1}. To be honest, here we should also mention the other works which are strongly criticizing the entropic force scenario. For instance, the author of~\\cite{r38} argued that there are some possible flaws in Verlinde's idea. In~\\cite{r39}, Culetu argued that the relativistic Unruh temperature cannot be associated with the bits on the screen in the form considered by Verlinde. In~\\cite{r40}, Hossenfelder argued that some additional assumptions made by Verlinde are unnecessary and there are some gaps in Verlinde's arguments. In~\\cite{r41}, Myung found that entropic force does not always imply the Newtonian force law, and the connection between Newtonian cosmology and entropic force cannot be confirmed. In~\\cite{r42}, Li and Pang found that inflation is inconsistent with the entropic force scenario. In~\\cite{r43}, Lee argued that there are some inconsistencies in Verlinde's arguments from a classical point of view.  So far, we have briefly surveyed the current status of the works relevant to the entropic force scenario. It is fair to say that the entropic force scenario is still in controversy. On the other hand, there is no breakthrough on entropic force after Verlinde's proposal~\\cite{r1}. A deep insight is needed to understand the nature of gravity. In addition, further discussions on the entropic force scenario are also desirable. Only when more and more results on entropic force are available, one can say something conclusively at that time. To this end, we would like to contribute our effort and try to learn more about the entropic force scenario. In this work, we will consider a modified entropic force scenario proposed by Gao~\\cite{r20}, which has some interesting features. And then, we will constrain the modified entropic force scenario with the latest observational data.  In~\\cite{r20}, Gao noted that statistical thermodynamics reveals the equipartition law of energy does not hold in the very low temperatures. Instead, as is well known, the Debye model~\\cite{r21,r22} is very successful in explaining the experimental results when the temperatures are very low. Since the equipartition law of energy plays an important role in the derivation of entropic force, the entropic force should be modified for the very weak gravitational fields which correspond to very low temperatures. Especially, the large-scale universe is in such an extreme weak gravitational field, and hence the modified entropic force (MEF) makes sense in cosmology.  Following~\\cite{r20}, we briefly mention the key points of MEF model. Similar to the Debye model~\\cite{r21,r22} in thermodynamics, one can modify the equipartition law of energy in Eq.~(\\ref{eq6}) to \\be{eq11} E=\\frac{1}{2}N k_B T D(x)\\,, \\ee where $D(x)$ is the Debye function which is defined by \\be{eq12} D(x)=\\frac{3}{x^3}\\int_0^x\\frac{y^3}{e^y-1}\\,dy\\,, \\ee and $x$ is related to the temperature $T$ as \\be{eq13} x\\equiv\\frac{T_D}{T}\\,, \\ee in which $T_D$ is the Debye temperature. By definition, $x$ is positive. With the modified equipartition law of energy, namely Eq.~(\\ref{eq11}), similar to the original entropic force, one can easily to obtain~\\cite{r20} \\be{eq14} g=\\frac{GM}{r^2}\\frac{1}{D(x)}\\,, \\ee in which (nb. Eq.~(\\ref{eq10})) \\be{eq15} x=\\frac{T_D}{T}=\\frac{g_D}{g}\\,, \\ee where $g_D\\equiv(2\\pi\\,k_B\\,c/\\hbar)\\,T_D$ is the Debye acceleration. Actually, Eq.~(\\ref{eq14}) corresponds to the modified Newtonian law of gravity. In the limit of strong gravitational field, $g\\gg g_D$ and hence $x\\ll 1$, from Eq.~(\\ref{eq12}) it is easy to find that $D(x)\\to 1$ and the Newtonian gravity is recovered. On the other hand, in the limit of weak gravitational field, $g\\ll g_D$ and hence $x\\gg 1$, one can see that $D(x)\\to\\pi^4/(5x^3)$ and then $g\\propto 1/\\sqrt{r}$, which significantly deviates from the familiar inverse square law~\\cite{r20}. However, as argued in~\\cite{r20}, one need not to worry about the possibility of MEF against the experimental results of the inverse square law. Since these experiments testing the inverse square law were done on the Earth or in the solar system, which are actually in the strong gravitational fields, we have $x\\ll 1$ and $D(x)\\to 1$, therefore the deviation from the inverse square law are extremely tiny. The significant deviation from the inverse square law can only occur in the very large scale in the universe where the gravitational fields are very weak, and hence it can escape the detection of these experiments testing the inverse square law. Of course, this argument relies on a small $g_D$. We will justify it later in the present work.  Next, we turn to the cosmological issues in the MEF Model. For convenience, we set the units $k_B=c=\\hbar=1$ hereafter. Using the derivation method in~\\cite{r5,r6}, one can find that the modified Raychaudhuri equation is given by~\\cite{r20} \\be{eq16} 4\\pi G\\left(\\rho+p\\right)=-\\left(\\dot{H}-\\frac{K}{a^2}\\right) \\left[-2D(x)+\\frac{3x}{e^x-1}\\right]\\,, \\ee where $\\rho$ and $p$ are the total energy density and total pressure of cosmic fluids, respectively; $K$ is the spatial curvature of the universe; $H\\equiv\\dot{a}/a$ is the Hubble parameter; $a=(1+z)^{-1}$ is the scale factor (we have set $a_0=1$); $z$ is the redshift; a dot denotes the derivatives with respect to cosmic time $t$; the subscript ``0'' indicates the present value of the corresponding quantity. Taking into account the Hawking temperature $T$ for the universe~\\cite{r23} \\be{eq17} T=\\frac{H}{\\,2\\pi}\\,, \\ee from Eqs.~(\\ref{eq13}) and (\\ref{eq15}), it is easy to see that \\be{eq18} x=\\frac{H_D}{H}\\,, \\ee where $H_D=g_D$. On the other hand, the energy conservation equation still holds in the MEF model, namely \\be{eq19} \\dot{\\rho}+3H\\left(\\rho+p\\right)=0\\,. \\ee From Eqs.~(\\ref{eq16}) and (\\ref{eq19}), one can derive the corresponding Friedmann equation. It is anticipated that Friedmann equation is also modified, $H^2\\not=(8\\pi G\\rho)/3$, due to the correction term $-2D(x)+3x/(e^x-1)$ in Eq.~(\\ref{eq16}). The MEF model is in fact a modified gravity model. Gao showed that the MEF model can describe the accelerating universe without dark energy. We refer to the original paper~\\cite{r20} for details.  Since Gao has not considered the constraints on the MEF model in~\\cite{r20}, we will try to obtain the cosmological constraints with the latest observational data in the next section. Further, we will discuss some relevant issues of the MEF model in Sec.~\\ref{sec3}. Finally, we give the brief conclusion and some meaningful remarks in Sec.~\\ref{sec4}.    \\begin{center} \\begin{figure}[tbhp] \\centering \\includegraphics[width=0.5\\textwidth]{sn.eps} \\caption{\\label{fig1} The $68.3\\%$ and $95.4\\%$ confidence level contours in the $\\Omega_{m0}-\\zeta$ parameter space. The best-fit parameters are also indicated by a black solid point. This result is obtained by using the data of 557 Union2 SNIa alone.} \\end{figure} \\end{center}   \\vspace{-10mm} % ",
        "conclusions": "\\label{sec4}  In summary, we considered the cosmological constraints on the MEF model in this work, by using the observational data of 557 Union2 SNIa, the distance parameter $A$ from LSS, the shift parameter $R$ from CMB, and 59 Hymnium GRBs. We found that the key parameter $\\zeta$ in MEF model has been limited in a narrow range $0\\leq\\zeta\\,\\lsim\\, 0.2$ within $95.4\\%$ confidence level. By using the important result given in Eq.~(\\ref{eq28}), we have clearly shown the key point to understand the reason for accelerating the universe without dark energy in the MEF model. We showed the MEF model reduces to $\\Lambda$CDM model when $\\zeta\\ll 1$. However, the best-fit $\\zeta$ for the observations SNIa+LSS+CMB and SNIa+LSS+CMB+GRBs significantly deviates from zero. This indicates the new feature of MEF model different from $\\Lambda$CDM model. We have justified the approximate solution $E(z)$ given in Eq.~(\\ref{eq28}). We plotted $E(z)$ and $q(z)$ as functions of redshift $z$, and clearly showed that the universe can be accelerated in late time without dark energy. Finally, we have shown that MEF can avoid the conflict with the experiments testing the inverse square law.  After all, some remarks are in order. In the MEF model, as shown in this work, the universe can be accelerated without dark energy. The only component is dust matter. The MEF model is in fact a modified gravity model, similar to the $f(R)$ models and the braneworld models. The MEF model can be degenerate to $\\Lambda$CDM model, but it has not an explicit cosmological constant in the model. This is a big advantage in fact, beyond some $f(R)$ and braneworld models in this sense.  Secondly, we point out the possibility to extend the original MEF model. In principle, it is not necessary to restrict the energy component in the universe to be dust matter only. The universe can contain other components, such as, dark energy. For example, we can consider a universe containing both dust matter and dark energy whose equation-of-state parameter (EoS) $w_X$ is a constant. In this case, the total energy density $\\rho=\\rho_m+\\rho_X=\\rho_{m0}\\,a^{-3}+\\rho_{X0}\\,a^{-3(1+w_X)}$. Substituting into Eq.~(\\ref{eq20}), we have \\be{eq37} \\left[-2D\\left(\\frac{\\zeta}{E}\\right)+\\frac{3\\,\\zeta/E} {e^{\\,\\zeta/E}-1}\\right]\\cdot 2E\\,\\frac{dE}{dz}= 3\\Omega_{m0}\\,(1+z)^2+ 3\\left(1+w_X\\right)\\Omega_{X0}(1+z)^{3w_X+2}\\,, \\ee where $\\Omega_{X0}\\equiv(8\\pi G\\rho_{X0})/(3H_0^2)$. Note that $\\Omega_{m0}+\\Omega_{X0}\\not=1$, since Friedmann equation has been modified in the MEF model. In principle, one can numerically solve Eq.~(\\ref{eq37}) to get $E(z)$. Similar to the original MEF model, we find that an approximation of the exact differential equation~(\\ref{eq37}) is given by \\be{eq38} \\left(1-\\frac{3}{4}\\frac{\\zeta}{E}\\right)\\cdot 2E\\,dE= \\Omega_{m0}\\,da^{-3}+\\Omega_{X0}\\,da^{-3(1+w_X)}\\,. \\ee Integrating Eq.~(\\ref{eq38}), we have \\be{eq39} E^2-\\frac{3}{2}\\,\\zeta E=\\Omega_{m0}\\,a^{-3} +\\Omega_{X0}\\,a^{-3(1+w_X)}+const., \\ee where $const.$ is the integral constant, which can be determined by requiring $E(z=0)=1$. Finally, we find that \\be{eq40} E^2-\\frac{3}{2}\\,\\zeta E=\\Omega_{m0}\\,(1+z)^3+ \\Omega_{X0}\\,(1+z)^{3(1+w_X)} +\\left(1-\\frac{3}{2}\\,\\zeta-\\Omega_{m0}-\\Omega_{X0}\\right), \\ee which is a quadratic equation of $E$ in fact. Noting that $E$ is positive, we solve Eq.~(\\ref{eq40}) to get \\be{eq41} E(z)=\\frac{3}{4}\\,\\zeta+\\frac{1}{2}\\left\\{\\frac{9}{4}\\,\\zeta^2 +4\\left[\\Omega_{m0}\\,(1+z)^3+\\Omega_{X0}\\,(1+z)^{3(1+w_X)} +\\left(1-\\frac{3}{2}\\,\\zeta-\\Omega_{m0}-\\Omega_{X0}\\right) \\right]\\right\\}^{1/2}. \\ee In fact, this is just a simple example. One can include any type of dark energy, for instance, the CPL dark energy whose EoS is given by $w_{de}=w_0+w_a(1-a)$, quintessence, phantom, k-essence, hessence, (generalized) Chaplygin gas, holographic/agegraphic dark energy, vector-like dark energy, spinor dark energy, and so on. Therefore, we would like to give the more general formulae. In this case, the total energy density $\\rho=\\rho_m+\\rho_{de}=\\rho_{m0}\\,a^{-3}+\\rho_{de,0}\\,f(a)$, where $f(a)$ can be any function of $a$ which satisfies $f(a=1)=1$. Substituting into Eq.~(\\ref{eq20}), we obtain \\be{eq42} \\left[-2D\\left(\\frac{\\zeta}{E}\\right)+\\frac{3\\,\\zeta/E} {e^{\\,\\zeta/E}-1}\\right]\\cdot 2E\\,\\frac{dE}{dz}= 3\\Omega_{m0}\\,(1+z)^2- \\Omega_{de,0}\\,(1+z)^{-2}f^\\prime\\,, \\ee where $f^\\prime\\equiv df/da$, and $\\Omega_{de,0}\\equiv(8\\pi G\\rho_{de,0})/(3H_0^2)$. Note again that $\\Omega_{m0}+\\Omega_{de,0}\\not=1$, since Friedmann equation has been modified in the MEF model. Eq.~(\\ref{eq42}) is the exact differential equation, which can be used to find the exact $E(z)$ numerically. Also, we give the corresponding approximate solution as \\be{eq43} E(z)=\\frac{3}{4}\\,\\zeta+\\frac{1}{2}\\left\\{\\frac{9}{4}\\,\\zeta^2 +4\\left[\\Omega_{m0}\\,(1+z)^3+ \\Omega_{de,0}\\,f\\left(\\frac{1}{1+z}\\right) +\\left(1-\\frac{3}{2}\\,\\zeta-\\Omega_{m0}-\\Omega_{de,0}\\right) \\right]\\right\\}^{1/2}. \\ee Similarly, if one need to add other component, such as radiation, it is not a hard work. In fact, we can give the most general formulae. In this case, the total energy density $\\rho=\\rho_0\\,f(a)$, where $f(a)$ can be any function of $a$ which satisfies $f(a=1)=1$. Substituting into Eq.~(\\ref{eq20}), we have \\be{eq44} \\left[-2D\\left(\\frac{\\zeta}{E}\\right)+\\frac{3\\,\\zeta/E} {e^{\\,\\zeta/E}-1}\\right]\\cdot 2E\\,\\frac{dE}{dz}= -\\Omega_0\\,(1+z)^{-2}f^\\prime\\,, \\ee where $\\Omega_0\\equiv(8\\pi G\\rho_0)/(3H_0^2)$. Note again that $\\Omega_0\\not=1$, since Friedmann equation has been modified in the MEF model, namely $H^2\\not=(8\\pi G\\rho)/3$. Eq.~(\\ref{eq44}) is the exact differential equation, which can be used to find the exact $E(z)$ numerically. Also, we give the corresponding approximate solution as \\be{eq45} E(z)=\\frac{3}{4}\\,\\zeta+\\frac{1}{2}\\left\\{\\frac{9}{4}\\,\\zeta^2 +4\\left[\\,\\Omega_0\\,f\\left(\\frac{1}{1+z}\\right) +\\left(1-\\frac{3}{2}\\,\\zeta-\\Omega_0\\right)\\right]\\right\\}^{1/2}. \\ee In fact, noting that $E=H/H_0$, Eq.~(\\ref{eq45}) can be regarded as the approximate modified Friedmann equation in the MEF model. If $\\zeta\\ll 1$, Eq.~(\\ref{eq45}) reduces to \\be{eq46} H^2=\\frac{8\\pi G}{3}\\,\\rho+\\Lambda_{\\rm eff}\\,, \\ee where $\\Lambda_{\\rm eff}=\\left(1-\\Omega_0\\right)H_0^2=const.$ is actually an effective cosmological constant. Therefore, in the most general case, we can clearly reveal the implicit root to accelerate the universe in the MEF model, regardless of the energy components in the universe. An effective cosmological constant is the intrinsic feature (we refer to e.g.~\\cite{r23} for a previous insight). The other exotic features of the MEF model could emerge only when $\\zeta$ significantly deviates from zero.  Thirdly, we would like to say some words on the understanding of the original entropic force model~\\cite{r1} and the modified entropic force model~\\cite{r20}. In fact, entropic force is just a new perspective to gravity, from the thermodynamical point of view. So, the original entropic force can only recover all results of the usual (Newton and Einstein) gravity. The only new thing is the reversed logic which might reveal the nature of gravity. In the original entropic force model~\\cite{r1}, using the fundamental assumptions Eqs.~(\\ref{eq1}), (\\ref{eq3}), (\\ref{eq5}) and~(\\ref{eq6}), Verlinde derived the Newton's law of gravitation Eq.~(\\ref{eq8}) for the (non-relativistic) Euclidean spacetime in section~3 of~\\cite{r1}, and also derived the Einstein gravitational equations for any (relativistic) curved spacetime in section~5 of~\\cite{r1}. On the other hand, the Friedmann equations were derived in~\\cite{r5,r6} for the Friedmann-Robertson-Walker (FRW) universe. There is {\\em no} any mixing here. We should mention that both the original entropic force~\\cite{r1} and the modified entropic force~\\cite{r20} cannot be understood in {\\em only} non-relativistic or relativistic cases. In fact, they are equivalent to gravity itself in all cases. As the usual understanding, the Newton's law of gravitation is just the approximation of Einstein gravitational equations in the (non-relativistic) small scale limit, whereas the Friedmann equations are just the special case of Einstein gravitational equations in the cosmic scale (homogeneous and isotropic spacetime). The situation is similar in the modified entropic force model. Using the fundamental assumptions Eqs.~(\\ref{eq1}), (\\ref{eq3}), (\\ref{eq5}) and~(\\ref{eq11}), in~\\cite{r20} Gao derived the Newton's law of gravitation Eq.~(\\ref{eq14}) for the (non-relativistic) Euclidean spacetime, and also derived the second Friedmann equation~(\\ref{eq16}) for the FRW universe. Note that the first Friedmann equation can be derived from the second Friedmann equation~(\\ref{eq16}) and the energy conservation equation~(\\ref{eq19}). On the other hand, following Verlinde's derivations in section~5 of~\\cite{r1}, one can derive the corresponding (modified) Einstein gravitational equations for any (relativistic) curved spacetime. In fact, this is just the lacked sector in the modified entropic force model. However, it is available in principle, although it has not been given in the literature. In the modified entropic force model, there is {\\em no} any mixing too. The modified Newton's law of gravitation Eq.~(\\ref{eq14}) is just the approximation of the (lacked but available in principle) modified Einstein gravitational equations in the (non-relativistic) small scale limit, whereas the modified second Friedmann equation~(\\ref{eq16}) is just the special case of the (lacked but available in principle) modified Einstein gravitational equations in the cosmic scale (homogeneous and isotropic spacetime).  Fourthly, we said that the MEF model is similar to $f(R)$-gravity or braneworld scenario. Notice that they are similar {\\em only} in the sense that the gravity has been modified in these models. Of course, both $f(R)$-gravity and braneworld scenario were derived from the known actions, whereas the action for MEF is still lacked in the literature. However, as mentioned above, following Verlinde's derivations in section~5 of~\\cite{r1}, in principle one can derive the corresponding (modified) Einstein gravitational equations for any (relativistic) curved spacetime. Once this lacked sector has been done, the explicit action is ready. Since the present work focuses on cosmology in the MEF model, we leave this task to future works.  Fifthly, as mentioned in this work, in the MEF model, there is {\\em no} dark energy in fact. The universe is matter-dominated always. The expansion of our universe is accelerated due to the fact that gravity has been modified. In the non-relativistic case, there is no dark energy too, but gravity is also modified. However, as mentioned in this work, this modification to Newtonian gravity is negligible on the Earth or in the solar system. In the larger scale, the modified gravity is described by Eq.~(\\ref{eq14}). As shown in~\\cite{r37}, the Debye entropic force can be an alternative to the modified Newtonian dynamics (MOND) to explain the rotational velocity curves of spiral galaxies. In fact, the MEF model~\\cite{r20} and the Debye entropic force model~\\cite{r37} are very similar. So, it is anticipated that the ``non-relativistic cosmology'' of the MEF model could be an alternative to dark matter, which is usually invoked to explain the rotational velocity curves of spiral galaxies.  Finally, we admit that the entropic force proposed by Verlinde is based on several unproved hypotheses, and it is still controversial in the physical community. On the other hand, the Debye model in the thermodynamics has not been used in the gravity theory previously. However, in the history, many great theories also appeared controversially in their beginning. Therefore, we consider that it is better to keep an open mind to these speculative attempts."
    },
    "1005/1005.4326_arXiv.txt": {
        "abstract": "We study the energy released from phase-transition induced collapse of neutron stars, which results in large amplitude stellar oscillations. To model this process we use a Newtonian hydrodynamic code, with a high resolution shock-capturing scheme. The physical process considered is a sudden phase transition from normal nuclear matter to a mixed phase of quark and nuclear matter. We show that both the temperature and the density at the neutrinosphere oscillate with time. However, they are nearly 180$^{\\circ}$ out of phase. Consequently, extremely intense, pulsating neutrino/antineutrino and leptonic pair fluxes will be emitted. During this stage several mass ejecta can be ejected from the stellar surface by the neutrinos and antineutrinos. These ejecta can be further accelerated to relativistic speeds by the electron/positron pairs, created by the neutrino and antineutrino annihilation outside the stellar surface. We suggest that this process may be a possible mechanism for short Gamma-Ray Bursts. ",
        "introduction": "Recently, by using simulations performed with the Newtonian numerical code introduced in \\cite{Lin06}, it was shown that the resulting quark star, produced by the phase transition induced collapse of a neutron star, will undergo a series of oscillations.  The collapse process with a conformally flat approximation to general relativity was also simulated in \\cite{Ab08}. The works of \\cite{Lin06,Ab08} focus on the gravitational wave signals emitted by the collapse process.  It is the purpose of the present paper to consider another important implication of this result, namely, the effect of the oscillations of the newly formed quark star on the neutrino emission. The oscillations can enhance the neutrino emission rate in a pulsating manner, and the neutrinos are emitted in a much shorter time scale. Moreover, through the process of neutrino-antineutrino annihilation, a large amount of electron-positron pairs is also produced. Therefore the neutron-quark phase transition in compact objects may be the energy source of GRBs \\cite{Ch09}. Such a model can also explain the lack of detection of a neutron star or pulsar formed in the SN 1987A \\cite{Cha09}, by assuming that the newly formed neutron star at the center of SN 1987A underwent a phase transition after the neutrino trapping timescale ($\\sim$ 10 s). Consequently,  the compact remnant of SN 1987A may be a strange quark star, which has a softer equation of state than that of neutron star matter. Such a phase transition can induce stellar collapse and result in large amplitude stellar oscillations. Extremely intense pulsating neutrino fluxes, with submillisecond period and with neutrino energy (greater than 30 MeV), can be emitted because the oscillations of the temperature and density are out of phase almost 180 $^{\\circ}$. If this is true, the current X-ray emission from the compact remnant of SN 1987A will be lower than $10^{34}$ erg s$^{-1}$, and it should be a thermal bremsstrahlung spectrum for a bare strange star with a surface temperature of around $\\sim $ 107 K. ",
        "conclusions": "In this paper we have studied the possible consequences of the phase-induced collapse of neutron stars to strange stars. We have found that both the density and the temperature inside the star will oscillate with the same period, but almost 180$^{\\circ}$ out of phase, which will result in the emission of intense pulsating neutrinos and  pairs. We want to point out that the intense pulse neutrino luminosity can be maintained due to the oscillatory fluid motion, which can carry thermal energy directly from the stellar core to the surface. This process can replenish the energy loss of neutrino emission much quicker than the neutrino diffusion process. A large fraction of the neutrino energy, roughly (1-1/$e$), will be absorbed by the matter very near the stellar surface. When this amount of energy exceeds the gravitational binding energy, some mass near the stellar surface will be ejected, and this mass will be further accelerated by absorbing pairs created from the neutrino and antineutrino annihilation processes outside the star. Although matter will be ejected periodically, each ejecta can have different masses and Lorentz factors, and therefore the intrinsic period could not be observed. We suggest that the collisions among these ejecta may produce short GRBs. Our numerical simulations are simulating a spherically symmetric and non-rotating  collapsing stellar object, and they also do not contain magnetic fields. Therefore the radiation emission produced in this model is isotropic. However, a realistic neutron star should have finite angular momentum and strong magnetic field, and hence these two factors could produce asymmetric mass ejection. This effect will be considered in future work.  The phase-transition from a neutron star to a strange star was simulated in \\cite{FrWo98}, with the conclusion that this process is most likely not a gamma-ray burst mechanism. They mimic the phase-transition by the arbitrary motion of a piston deep within the star, and they have found that the mechanic wave will eject $\\sim 10^{-2}M_{\\odot}$ baryons, which causes the baryon contamination for the gamma-ray bursts. In our simulations, we assume a sudden change of equation of state to mimic the phase-transition, and we use the Newtonian hydrodynamic code to study the response of the stellar interior after such a sudden change of the EOS. In our simulations we find that the mass ejection by the motion of the fluid is very small. We estimate that the major mass ejection would result from the heating of neutrinos and pairs on the stellar crust, which is not modeled in the simulations. Our total energy output and total mass in ejecta are close to that of \\cite{FrWo98}. However, the neutrino energy injection is pulsating, and hence the mass ejection is also pulsating. The mass of individual ejecta range from $\\sim 10^{-9}$ to $\\sim 10^{-4}M_{\\odot}$, with output energy in the range of $10^{48}$ to $10^{50}$ ergs. Therefore, some ejecta cannot be relativistic, and they cannot contribute to GRBs. However, there are still many relativistic ejecta in each simulation model, which can have Lorentz factors $>$100, and with a total energy of $\\sim 10^{50}$ --- $10^{51}$ ergs. This could be a possible mechanism for short GRBs.    \\bigskip KSC and TH are supported by the GRF Grants of the Government of the Hong Kong SAR under HKU7013/06P and HKU7025/07P respectively."
    },
    "1005/1005.4332_arXiv.txt": {
        "abstract": "{ {The stable propagation of jets in FRII sources is remarkable if one takes into account that large-scale jets are subjected to potentially highly disruptive three-dimensional (3D) Kelvin-Helmholtz instabilities.} {Numerical simulations can address this problem and help clarify the causes of this remarkable stability. Following previous studies of the stability of relativistic flows in two dimensions (2D), it is our aim to test and extend the conclusions of such works to three dimensions.} {We present numerical simulations for the study of the stability properties of 3D, sheared, relativistic flows. This work uses a fully parallelized code (\\textit{Ratpenat}) that solves equations of relativistic hydrodynamics in 3D.} {The results of the present simulations confirm those in 2D. We conclude that the growth of resonant modes in sheared relativistic flows could be important in explaining the long-term collimation of extragalactic jets.} {} ",
        "introduction": "\\label{intro}  Extragalactic jets from AGN present a morphological dichotomy between FRI and FRII type jets \\citep{fr74}. The former \\citep[e.g., 3C~31,][]{lb02} show a disrupted structure at kiloparsec scales, whereas the latter \\citep[e.g., Cyg A,][]{cb96} are highly collimated. Although the origin of this dichotomy is a complex combination of several intrinsic (e.g., jet power) and external (i.e., environmental) factors, the stable propagation of jets in FRII sources is remarkable considering that large-scale jets are subjected to potentially highly disruptive three-dimensional (3D) modes of the Kelvin-Helmholtz (KH) instability. Any perturbation may couple to an instability and grow in amplitude, becoming potentially disruptive \\citep{pe+05}, hence the importance of studying their development and growth from the linear to the nonlinear regime in terms of the flow parameters.  Jet stability in the relativistic regime has been thoroughly studied in different scenarios and parameter ranges \\citep[see][for a recent review on the state-of-art]{ha06}, from both the analytical and numerical perspectives. The linear analysis of KH instability in relativistic flows started with the work of \\cite{ts76} and \\cite{bp76}, who derived and solved a dispersion relation for a single plane boundary between a relativistic flow and the ambient medium. Next, \\cite{ftz78} and \\cite{ha79} examined properties of the KH instability in relativistic cylindrical jets. The effects of a shear layer have been examined by \\cite{ur02}. More recently, \\cite{ha08} has studied the stability of magnetized relativistic jets with poloidal magnetic fields.  An extension of the linear stability analysis in the relativistic case to the weakly nonlinear regime has been performed by \\cite{hz95,hz97} and led to the conclusion that KH instability saturates at finite amplitudes because of various nonlinear effects. The most significant effect results from the relativistic character of the jet flow, namely from the velocity perturbation not exceeding the speed of light in the reference frame of the jet. This result was confirmed by numerical simulations \\citep{pe+04a,pe+04b}.  In the purely nonlinear regime, \\cite{ma97}, \\cite{ha98} and \\cite{ro99} demonstrate that jets with high Lorentz factors and high internal energy are influenced very weakly by the Kelvin-Helmholtz instability, contrary to the cases with lower Lorentz factors and lower internal energies. The former do not develop modes of KH instability predicted by the linear theory, which is interpreted as the result of a lack of appropriate perturbations (triggered by the backflow in this kind of simulations) generating the instability in the system, because in the limit of high internal energies of the jet matter the Kelvin-Helmholtz instability is expected to develop with the highest growth rate.  A combination of linear and nonlinear studies of jet stability in the relativistic case was applied for the first time by \\cite{ha98} in the case of axisymmetric, cylindrical jets and then extended to the 3D case by \\cite{ha01} in the spatial approach and by \\cite{pe+04a,pe+04b} in the temporal approach, for 2D slab and cilyndrical jets, including the case of sheared flows in \\cite{pe+05}. Recently, \\cite{mi07} have extended this study to 3D relativistic magnetized jets, which may be stabilized (with respect to current-driven modes) by a sheath layer surrounding the jet core.  Focusing on the combination of linear and nonlinear studies, \\cite{pe+04a,pe+04b} studied the growth of single symmetric modes in relativistic jets, sweeping different parameters, such as jet temperature, Lorentz factor, and density ratio, with the external medium. The linear work confirmed that KH modes grow faster in slower and hotter flows than in cold and fast ones. In Perucho, Mart\\'{\\i} \\& Hanasz (2005, Paper I from now on), this work was extended to sheared jets to study the growth of a mixture of symmetric and antisymmetric modes in this case. In \\cite{pe+04b} and Paper~I, the authors determined the stability properties of relativistic flows for a wide range of parameters in density contrast ($10^{-4}-0.1$), Lorentz factor ($5-20$) and specific internal energy ($0.08-76.5\\,\\rm{c^2}$). It was shown that the inertia of the jet, in terms of the density contrast with the external medium, the relativistic Mach number and the Lorentz factor are the  parameters that determine the stability of jets in the nonlinear regime, as a relativistic counterpart to the work of \\cite{bm94} in the classical case. Moreover, in Paper I and \\cite{pe+07} (Paper II, from now on), it was shown that the growth of high-order Kelvin-Helmholtz modes developing in the shearing layer, hereafter referred to as resonant modes, could dramatically change the nonlinear evolution of the flow in the case of cold and fast jets.  The results could be summarized as follows. a) Cold and slow jets (small relativistic Mach number and Lorentz factor) are unstable with the growth of the instability and are disrupted after the generation of a shock front crossing the boundary between the jet and the ambient medium, which results from the development of long wavelength KH instability modes (UST1 jets). b) Hot and slow jets (intermediate values of relativistic Mach numbers and Lorentz factor) are also unstable, but disrupted and mixed in a continuous way by the growth of the mixing layer down to the jet axis (UST2 jets). c) Faster (high values of Mach number and/or Lorentz factor) jets develop short-wavelength, high-order modes that grow in the shear layer and saturate the growth of the instability without loss of either collimation or mixing, and generate a hot shear layer around the core of the jet (ST jets). d) Lighter jets are more unstable than denser ones (UST1).  In this work, we extend the previous studies to 3D and show the existence of a new mechanism to prevent jet decollimation based on the development of resonant KH perturbations. With the aforementioned experience in mind, and taking into account that 3D simulations are highly demanding in terms of computational resources, we present here three numerical simulations using parameters that are representative of the different stability regions found in Paper I, but now including helical and elliptic modes. It is our aim to test whether the same conclusions are valid and, thus, if resonant modes provide a general stabilising mechanism in the absence of strong magnetic fields.  The paper is organized as follows. Section~\\ref{code} is devoted to presenting of a new 3D-RHD code for these simulations and the numerical setup used. Section~\\ref{res} contains the results of the linear analysis for the parameters considered and the simulations. In Section~\\ref{disc}, the results are discussed and the conclusions of this work are presented in Section~\\ref{conc}. ",
        "conclusions": "\\label{conc} In this work, we presented the first solutions of the linear problem for 3D sheared relativistic flows. Following this result, we presented numerical simulations of the growth of KH instability modes from the linear to the nonlinear regimes, for three selected relativistic jet models. These numerical simulations were performed with the new high-resolution shock-capturing 3D RHD code \\textit{Ratpenat}, which was parallelized using a hybrid scheme with OMP and MPI processes. The simulations of jets in the typical stability regions of the space of parameters (Paper~I) have confirmed previous results obtained in the 2D approximation. We proved that under certain conditions, the appearance of the shear-layer resonant modes of KH instability has important consequences for the stability of relativistic jets. In particular, the implications for the collimation of FRII jets were discussed. We also show, as a by-product of our simulations, that the development of different modes with a variety of wavelengths may induce rich transversal structure in jets, with shorter modes showing up in the central regions of the jet and longer ones dominating the structure in the mixing layer. In this respect, we confirm the results obtained in \\cite{pe+06}.  Future work along this line will consider the extension of our study in both the linear and nonlinear regime by means of simulations with higher resolution to a wider sample of jet parameters, including magnetic fields and a more realistic equation of state."
    },
    "1005/1005.4618_arXiv.txt": {
        "abstract": "We present additional evidence that dust is really forming along the red giant branch (RGB) of 47~Tuc at luminosities ranging from above the horizontal branch to the RGB-tip \\citep{ori07}. The presence of dust had been inferred from an infrared excess in the $(K-8)$ color, with $K$ measured from high spatial resolution ground based near-IR photometry and ``8'' referring to Spitzer-IRAC 8\\micron\\ photometry. We show how $(K-8)$ is a far more sensitive diagnostic for detecting tiny circumstellar envelopes around warm giants than colors using only the Spitzer-IRAC bands, for example the $(3.6-8)$ color used by \\citet{boy10}. In addition, we also show high resolution HST-ACS $I$ band images of the giant stars which have $(K-8)$ color excess. These images clearly demonstrate that \\citet{boy10} statement that our detections of color excess associated with stars below the RGB-tip arise from blends and artefacts is simply not valid. ",
        "introduction": "\\label{intro}  In \\citet{ori07} we presented estimated mass loss rates for first ascent red giant branch (RGB) stars in the globular cluster 47~Tuc. These were based on near and mid-infrared photometry obtained from a Spitzer-IRAC (3.6, 4.5, 5.6 \\& 8\\,\\micron) survey and ground-based high resolution $JHK$ observations.  We found about 100 giants with ($K-8$) color excess that we attributed to the presence of dusty circumstellar envelopes.  These candidate dusty stars were mainly found in the inner 2\\arcmin\\ (in radius) and had luminosities ranging from above the level of the Horizontal Branch (HB) to the Red Giant Branch (RGB)-tip.  For a given luminosity only a fraction of stars exhibited this color excess, and this fraction increases toward the tip of the RGB.  From the color excess we derived a mass loss rate and found a shallower dependence on luminosity than that expected from Reimers \\citep{rei75a,rei75b} formula. After correcting the observed frequency of dusty envelopes for incompleteness, we derived an average duty cycle which could be used along with the observed mass loss rates and evolutionary times to compute the total amount of mass lost on the RGB.  Recently, \\citet{boy10}, using the $(3.6-8)$ Spitzer-IRAC color as their main diagnostic tool, found evidence for a dust excess only in asymptotic giant branch stars and possibly a few giants near the RGB tip. On the basis of this finding, they concluded \\citep[see also ][]{boy08,boy09} that our candidate dusty giants below $M_{\\rm bol} = -2.5$ were spurious, mainly blends and/or artefacts.  We emphasize that {\\it i}) this conclusion was not based on a direct star-to-star check, {\\it ii}) their analysis employed a different diagnostic, and it was not optimized to cover the innermost region of the cluster sampled by our work.  The AKARI analysis of 47~Tuc presented by \\citet{ita07} also showed dust excess only near the tip, but it had not the spatial resolution to properly investigate the stellar population in central regions of globular clusters.  We ourselves using ISOCAM found circumstellar dust excess only near the RGB tip \\citep{ori02}, again because of the lower spatial resolution of ISOCAM compared to IRAC.  In this Letter we respond to the \\citet{boy10} criticism. We demonstrate the importance of an optimum sampling of the innermost region of the cluster where most of the red giants, along with the rest of the stellar population, are found. We show that $(K-8)$ is a much better diagnostic of dusty giants than $(3.6-8)$. Finally, using a high resolution HST-ACS $I$ band image, we clearly demonstrate that our detected $(K-8)$ color excess in giant stars below the RGB tip is not an effect of blending and/or an artefact. ",
        "conclusions": "The $(K-8)$ color excess detected by \\citet{ori07} in a fraction of RGB stars down to $\\mbol \\sim 0$ within the central 2\\arcmin\\ of 47~Tuc is real. It is not an artefact, since blends and other possible spurious effects have been properly accounted for by cross-correlating the Spitzer sources with high spatial resolution optical and near infrared catalogs of stellar counterparts.  The fact that \\citet{boy10} did not find as many dusty stars in the central 2\\arcmin\\ is mainly due to the use of a different diagnostic, namely the $(3.6-8)$ color, which is only effective for detecting dust excess in cooler (hence more luminous) giants. Further, their photometric analysis is not optimized in the central, densest region. \\citet{boy10} find only few dusty giants in the outer regions, fully consistent with \\citet{ori07} and the additional CMDs shown in Figure 2.  As discussed in Section 4 of \\citet{ori07}, mass loss rates and duty cycles decrease with decreasing stellar luminosity.  For a given luminosity, the estimated rates are higher (between a factor of 2 near the tip up to 2 orders of magnitude down to $\\mbol \\sim 0$ than those predicted by the Reimers law, as also noted by \\citet{boy10}. The shallower slope we found has been also suggested by \\citet{mes09}, from chromospheric line wind diagnostics of giants in metal poor clusters.  The Reimers result was obtained from Population~I objects, and one might anticipate some differences.  Our higher mass loss rate is not a major problem, given that both the Reimers law and our results contain free parameters, {\\it i}i) the Reimers efficiency $\\eta$, and {\\it ii}) the gas to dust ratio and expansion velocity. In either case the free parameters must be set by indirect observational constraints on the total mass lost during the RGB evolution like the HB morphology.  The total mass loss occurs predominantly in the upper $\\sim2$ magnitudes near the RGB tip. The contribution of the low luminosity giants ($\\mbol > -1$) is small ($<20$\\%) to negligible, and within the estimated uncertainty.  The importance of our observation of mass loss in less luminous stars is not its impact on total mass lost, but rather the clue it gives us to the physics of mass loss. Likewise, the episodic nature of mass loss in non-variable stars tells us something. In both cases, it seems quite clear that the underlying driver of the mass loss is {\\em not} radiation pressure on dust. Another important outcome of our larger project will be an investigation of the differential mass loss among clusters with different metallicities and HB morphologies.  This has been the goal of our Spitzer survey and the complementary ground based observations, and the results will appear in forthcoming papers."
    },
    "1005/1005.1451_arXiv.txt": {
        "abstract": "Using the state-of-the-art cosmological hydrodynamic simulations of the standard cold dark matter model with star formation feedback strength normalized to match the observed star formation history of the universe at $z=0-6$, we compute the metal enrichment history of the intergalactic medium (IGM). Overall we show that galactic superwind (GSW) feedback from star formation can transport metals to the IGM and that the properties of simulated metal absorbers match current observations. The distance of influence of GSW from galaxies is typically limited to about $\\le 0.5$Mpc and within regions of overdensity $\\delta \\ge 10$. Most \\civ and \\ovi absorbers are located within shocked regions of elevated temperature ($T\\ge 2\\times 10^4$K), overdensity ($\\delta \\ge 10$), and metallicity ($[Z/\\zsun]=[-2.5,-0.5]$), enclosed by double shocks propagating outward. \\ovi absorbers have typically higher metallicity, lower density and higher temperature than \\civ absorbers. For \\ovi absorbers collisional ionization dominates over the entire redshift range $z=0-6$, whereas for \\civ absorbers the transition occurs at moderate redshift $z\\sim 3$ from collisionally dominated to photoionization dominated. We find that the observed column density distributions for \\civ and \\ovi in the range $\\log N {\\rm cm}^2=12-15$ are reasonably reproduced by the simulations. The evolution of mass densities contained in \\civ and \\ovi lines, $\\omegaciv$ and $\\omegaovi$, is also in good agreement with observations, which shows a near constancy at low redshifts and an exponential drop beyond redshift $z=3-4$. For both \\civ and \\ovi\\, most absorbers are transient and the amount of metals probed by \\civ and \\ovi lines of column $\\log N {\\rm cm}^2=12-15$ is only $\\sim 2\\%$ of total metal density at any epoch. While gravitational shocks from large-scale structure formation dominate the energy budget ($80-90\\%$) for turning about 50\\% of IGM to the warm-hot intergalactic medium (WHIM) by $z=0$, GSW feedback shocks are energetically dominant over gravitational shocks at $z\\ge 1-2$. Most of the so-called ``missing metals\" at $z=2-3$ are hidden in a warm-hot ($T=10^{4.5-7}$K) gaseous phase, heated up by GSW feedback shocks. Their mass distribution is broadly peaked at $\\delta=1-10$ in the IGM, outside virialized halos. Approximately $(37,46,10,7)\\%$ of the total metals at $z=0$ are in (stars, WHIM, X-ray gas, cold gas); the distribution stands at $(23,57,2,18)\\%$ and $(14,51,4,31)\\%$ at $z=2$ and $z=4$, respectively. ",
        "introduction": "One of the pillars of the Big Bang theory is its successful prediction of a primordial baryonic matter composition, made up of nearly one hundred percent hydrogen and helium with a trace amount of a few other light elements \\citep[e.g.,][]{1998Schramm,2001Burles}. The metals, nucleosynthesized in stars later, are found almost everywhere in the observable IGM, ranging from the metal-rich intracluster medium \\citep[e.g.,][]{1997Mushotzky} to moderately enriched damped Lyman systems \\citep[e.g.,][]{1997Pettini, 2003Prochaska} to low metallicity Lyman alpha clouds  \\citep[e.g.,][]{2003Schaye}. When and where were the metals made and why are they distributed as observed? We address this fundamental question in the context of the standard cold dark matter cosmological model \\citep[][]{2009Komatsu} using latest simulations. Our previous simulations \\citep[][]{1999bCen, 2005Cen} provided some of the earlier attempts to address this question with measured successes. In this investigation we use substantially better simulations to provide significantly more constrained treatment of the feedback processes from star formation (SF) that drive energy and metals from supernovae into the IGM through galactic winds \\citep[e.g.,][]{1999bCen, 2001Aguirre, 2002bTheuns, 2003Adelberger, 2003Springel}.   Metal-line absorption systems in QSO spectra are the primary probes of the metal enrichment of the IGM as well as in the vicinities of galaxies \\citep[e.g.,][]{1969Bahcall}. The most widely used metal lines include \\ion{Mg}{2} $\\lambda \\lambda$2796, 2803 doublet \\citep[e.g.,][]{1992Steidel}, \\ion{C}{4} $\\lambda \\lambda$1548, 1550 doublet \\citep[e.g.,][]{1982Young}, and \\ion{O}{6} $\\lambda \\lambda$1032, 1038 doublet \\citep[e.g.,][]{2002Simcoe}. We here focus on the \\civ and \\ovi absorption lines and the global evolution of metals in the IGM. We will limit our current investigation to the observationally accessible redshift range of $z=0-6$, which in part is theoretically motivated simply because the theoretical uncertainties involving still earlier star formation are much larger. At $z=0$ the \\ovi line (together with \\ion{C}{7} and \\ion{O}{8} lines) provide vital information on the missing baryons \\citep[e.g.,][]{2003Mathur, 2008Tripp, 2008Danforth, 2009Nicastro}, predicted to exist in a Warm-Hot Intergalactic Medium (WHIM) \\citep[][]{1999bCen, 2001Dave}.  For a well understood sample of QSO absorption lines, one could derive the cosmological density contained in them \\citep[e.g.,][]{2009Cooksey}. Early investigations indicate that $\\omegaciv$ remains approximately constant in the redshift interval $z \\sim 1.5 - 4$ \\citep[][]{2001Songaila,2005Songaila,2003Boksenberg}. There have been recent efforts to extend the measurements of  $\\omegaciv$ to $z<1.5$ \\citep[][]{2009Cooksey} and to $z>5$ \\citep[][]{2006Simcoe,2006Ryanweber, 2009Ryanweber,2009Dodorico, 2009Becker}.  Observations in these redshift ranges have been difficult to carry out because \\civ transition moves to the UV at low redshift and to the IR band at high redshift. \\citet[][]{2009Dodorico} find evidence of a rise in the \\civ mass density for $z<2.5$. \\citet[][]{2006Simcoe} and \\citet[][]{2006Ryanweber} found evidence of \\civ density at $z\\sim 6$ being consistent with estimations at $z \\sim 2-4.5$. More recently, however, \\citet[][]{2009Becker} set upper limits for $\\omegaciv$ at $z\\sim 5.3$ and \\citet[][]{2009Ryanweber} observe a decline in intergalactic \\civ approaching $z=6$, which we will show are in good agreement with our simulations.  The ionization potential of \\ovi and the relatively high oxygen abundance are very favorable for production of \\ovi absorbers in the IGM \\citep[e.g.,][]{1983Norris, 1986Chaffee}. The rest wavelength of OVI ($1032,1037$\\AA) places it within the Ly-$\\alpha$ forest, which makes the identifications of these lines more complicated, although being a doublet helps significantly. At $z\\ge 2 $, however, \\ovi absorption can probe the metal content of the IGM in ways complementary to what is provided by \\civ lines. For example, the \\ovi lines can probe IGM that is hotter than that probed by the \\civ lines and can reach lower densities thank to higher abundance. There are now several observational studies at redshifts $z=2-3$ that describe the properties of \\ovi absorbers and attempt to estimate the \\ovi mass density, $\\omegaovi$ \\citep[][]{2002Carswell,2002Bergeron, 2004Simcoe, 2006Simcoe, 2008Frank, 2008Danforth, 2008Tripp, 2008Thomb}.   At $z\\sim 2-3$ there is a missing metals problem: only 10-20\\% of the metals produced by all stars formed earlier have been identified in stars of Lyman break galaxies (LBG), in damped Lyman alpha systems (DLAs) and $\\lya$ forest. The vast majority of the produced metals appear to be missing \\citep[e.g.,][]{1999bPettini}. The missing metals could be in hot gaseous halos of star-forming galaxies \\citep[][]{1999bPettini, 2005Ferrara}. We will show that most of the missing metals are in a warm-hot ($T=10^{4.5-7}$K) but diffuse IGM at $z=2-3$ of overdensities of $\\sim 10$ that are outside of halos.    The outline of this paper is as follows. In \\S 2 we detail our simulations and the procedure of normalizing the uncertain feedback processes from star formation. Results on the metal enrichment of the IGM are presented in \\S 3. In \\S 3.1 we give a full description of the properties of the \\civ and \\ovi lines at $z=0-6$, followed \\S 3.2 discussing \\civ and \\ovi absorbers as metals reservoirs. We devote \\S 3.3 to a general discussion of global distribution of metals, addressing several specific topics, including the metallicity of the moderate overdense regions at moderate redshift, the missing metals at $z\\sim 3$. Conclusions are given in \\S 4. ",
        "conclusions": "We have carried out the state-of-the-art cosmological hydrodynamic simulations of the standard cold dark matter model to investigate the process of metal enrichment of the intergalactic medium. Our simulations have substantially higher resolution than our previous simulations to address this problem. More importantly, we can now constrain the strength of the feedback process by matching the star formation history in our current simulation to the observed one in the range $z=0-6$. We find that our model reproduces the observed mean flux of the $\\lya$ forest and the mass density of \\civ and \\ovi absorbers. It is also in general consistent with observed physical properties of absorbers. This indicates that we can explain the metal enrichment of the IGM by considering star formation to be the main feedback mechanism, with no apparent need of significant contribution from AGN in terms of additional energy. We conclude from our results that:   (1) The overall star formation history depends rather sensitively on the feedback strength. This is likely due to GSW significantly reducing the concentration of cold gas around halos. Nevertheless, GSW do not significantly alter the overall large-scale filamentary baryonic structure that follows the cosmic web of dark matter distribution. While GSW could travel far into low density regions sometimes, the amount of energy and metals that are deposited in underdense regions is very small. Most of the GSW energy and metals remain in regions of overdensity $\\delta \\ge 10$, with the distance of influence of GSW from galaxies limited to about $\\le 0.5$Mpc. Metal bubbles blown by GSW coincide with temperatures bubbles, suggesting a tight coupling of energy and metal deposition, and they are terminated by shock fronts.  (2) Both \\civ and \\ovi absorbers are located in regions that have been swept by feedback shocks, of elevated temperature ($T\\ge 2\\times 10^4$K), density ($\\delta \\ge 10$) and metallicity ($[Z/\\zsun]=[-2.5,-0.5]$), demarcated by a double shock propagating outwards, with \\ovi absorbers typically having a higher metallicity than \\civ absorbers. Within these shocked regions, most of \\civ absorbers tend to arise from moderate density peaks that are troughs in temperature and are thus relatively quiescent. The \\ovi absorbers are from regions that are dynamically hotter near shock fronts. There is a trend for the population of \\civ and \\ovi absorbers to be more collisionally ionized at higher redshift; for \\ovi collisional ionization dominates over the entire redshift range $z=0-6$, whereas for \\civ the transition occurs at moderate redshift $z\\sim 3$ from collisionally dominated to photoionization dominated.  (3) The evolution of the mass density contained in \\civ and \\ovi lines, $\\omegaciv$ is in good agreement with observations, with both the latest observations and simulations of $\\omegaciv$ exhibiting an exponential drop beyond redshift $z=4$; $\\omegaciv$ drop exponentially beyond redshift $z=3$; the near constancy of $\\omegaciv$ at redshift $z=1-3$ does not reflect the evolution of the overall metal content in the IGM. In the case of $\\omegaovi$, we find a less good agreement between observations and out results. This might be in part due to cosmic variance.  (4) Most of \\civ and \\ovi absorbers, while clustered around galaxies, are transient and intergalactic in origin, produced by galactic superwinds in the process of transporting both energy and metals from galaxies into the IGM; the metal mass densities contained in \\civ and \\ovi lines in the range $\\log N {\\rm cm}^2=12-15$ each constitutes $\\sim 0.1\\%$ of total metal density at all redshifts; the amount of metals probed by \\civ and \\ovi lines in the range $\\log N {\\rm cm}^2=12-15$ is $\\sim 1\\%$ of the total metal density at all redshifts.   (5) While gravitational shocks from large-scale structure formation dominate the energy budget ($80-90\\%$) for turning about $50\\%$ of IGM to the warm-hot intergalactic medium (WHIM) by $z=0$, galactic superwind feedback shocks are energetically dominant over gravitational shocks at $z\\ge 1-2$.  (6) Most of the so-called ``missing metals\" at $z=2-3$ are hidden in a warm-hot gaseous phase ($T>3\\times 10^4$K) that is heated up by star formation feedback shocks. Their mass distribution is broadly peaked at overdensity $1-10$ in the IGM, outside virialized halos. Approximately $(37,46,10,7)\\%$ of the total metals at $z=0$ are in (stars, WHIM, X-ray gas, cold gas); the distribution stands at $(23,57,2,18)\\%$ and $(14,51,4,31)\\%$ at $z=2$ and $z=4$, respectively.   (7) The metallicity of the IGM with moderate overdensities ($1-10$) that are probed by the $\\lya$ forest shows a rapid increase with decreasing redshift. We show that velocity ``diffusion'' effect that arises from the peculiar velocities could enhance the ``apparent\" metallicity of the $\\lya$ forest clouds, as supported by our cross-correlation analysis. Tentatively, we suggest that this may reconcile, at least in part, the discrepancy between our simulations and observations at $z=2-4$ based on pixel optical depth (POD) method."
    },
    "1005/1005.0382_arXiv.txt": {
        "abstract": "{  Past and current X-ray mission allow us to observe only a fraction of the volume occupied by the ICM. After reviewing the state of the art of cluster outskirts observations we discuss some important constraints that should be met when designing an experiment to measure X-ray emission out to the virial radius. From what we can surmise \\wfxt ~ is already designed to meet most of the requirements and should have no major difficulty in accommodating the remaining few. ",
        "introduction": "Galaxy clusters form through the hierarchical accretion of cosmic matter. The end products of this process are virialized structures that feature, in the X-ray band, similar radial profiles of the surface brightness $S_{\\rm b}$ (e.g. Vikhlinin et al. 1999, Neumann 2005, Ettori \\& Balestra 2009) and of the plasma temperature $T_{\\rm gas}$ (e.g. Allen et al. 2001, Vikhlinin et al. 2005, Leccardi \\& Molendi 2008). Such measurements have definitely improved in recent years thanks to the arcsec resolution and large collecting area of the present X-ray satellites, like \\chandra\\ and \\xmm, but still remain difficult, in particular in the outskirts, because of the low surface brightness associated to these regions. Present observations provide routinely reasonable estimates of the gas density, $n_{\\rm gas}$, and temperature, $T_{\\rm gas}$, up to about $R_{2500}$ ($\\approx 0.3 R_{200}$; $R_{\\Delta}$ is defined as the radius of the sphere that encloses a mean mass density of $\\Delta$ times the critical density at the cluster's redshift; $R_{200}$ defines approximately the virialized region in galaxy clusters). Only few cases provide meaningful measurements at $R_{500}$ ($\\approx 0.7 R_{200}$) and beyond (e.g. Vikhlinin et al. 2005, Leccardi \\& Molendi 2008, Neumann 2005, Ettori \\& Balestra 2009). Consequently, more than two-thirds of the typical cluster volume, just where primordial gas is accreting and dark matter halo is forming, is still unknown for what concerns both its mass distribution and its thermodynamical properties. This poses a significant limitation in our ability to characterize the physical processes presiding over the formation and evolution of clusters and to use clusters as cosmological tools, as also outlined in the Scientific Justification for the \\wfxt\\ (Giacconi et al. 2009). Indeed the characterization of thermodynamic properties at large radii would allow us to provide constraints on the virialization process, while measures of the metal abundance would allow us to  gain insight on the enrichment processes occurring in clusters (e.g. Fabjan et al. 2010). Morever the X-ray emission at large radii could also be used to improve significantly measures of the gas and total gravitating masses thereby opening the way to a more accurate use of galaxy clusters as cosmological probes (e.g. Voit 2005).  In these proceedings, we take stock of the situation on cluster outskirts and suggest how to make progress. In Sect.~\\ref{sec: obs},  we provide an observational overview of currently available measures of cluster outer regions,  while in Sect.~\\ref{sec: map} we discuss some important constraints that should be met  when designing an experiment to measure X-ray emission out to the virial radius. In Sect.~\\ref{sec: future}, we present an overview of future missions which have cluster outskirts observations as one of their goals, our main results are recapitulated  in Sect.~\\ref{sec: summary}.   \\begin{figure*} \\begin{center} \\includegraphics[height=4.5cm]{eb09_fig2.ps} \\end{center} \\caption{\\small {\\bf From left to right}: Example of a surface brightness profile with the fitted background ({\\it horizontal dotted line}) and the radius $R_{200}$ ({\\it vertical dashed line}); the signal-to-noise profile evaluated as $S2N = (S_b - B)/ \\epsilon$, where the error $\\epsilon$ is the sum in quadrature of the Poissonian error in the radial counts and the uncertainties in the fitted background, $B$; the best-fit values of the slope of the surface brightness profile as a function of $r / R_{200}$. These values are estimated over 6 radial bins (thick horizontal solid line: the slope evaluated between $0.4 \\times R_{200}$ and $R_{S2N}$ with a minimum of 3 radial bins; dashed line: best-fit of $d \\ln (S_b) / d \\ln (r/R_{200})$ with the functional form $s_0 +s_1 \\ln (r/R_{200})$ over the radial range $0.1 \\times R_{200} - R_{S2N}$, with the best-fit parameters quoted in Table~3 of Ettori \\& Balestra 2009). } \\label{fig:s2n} \\end{figure*}  A Hubble constant of 70 $h_{70}$ km s$^{-1}$ Mpc$^{-1}$ in a flat universe with $\\Omega_{\\rm m}$ equals to 0.3 is assumed throughout this manuscript. ",
        "conclusions": "\\label{sec: summary} Past and current X-ray mission allow us to observe only a fraction of the volume occupied by the ICM. Indeed, typical measures of the surface brightness, temperature and metal abundance extend out to a fraction of the virial radius. The coming into operation of the second generation of medium energy X-ray telescopes at the turn of the millennium, has resulted in relatively modest improvements in our ability to characterize cluster outskirts. Even though recent results from \\suzaku\\ show some improvement, the  most sensitive instrument to low surface brightness to have flown thus far is quite possibly the \\swift\\ XRT which, ironically, never had cluster outer regions as one of its top scientific objectives.  The construction of an experiment capable of making measures out to $R_{200}$ is well within the reach of currently available technology. What is required is an experiment design that will minimize the background, both instrumental and cosmic, and maximizes the grasp, i.e. the product of effective area and FOV. Since cluster emission in the outskirts will be background dominated, instrument design and observational strategy should also allow for a meticulous characterization of the background. Detailed simulations based on realistic estimates of the different spectral components and of the precision with which the may be determined shows that an experiments such as the one we envisage will allow a solid characterization of cluster outskirts. From what we can surmise \\wfxt\\ is already designed to meet most of the requirements which are necessary to characterize cluster outskirts, and should have no major difficulty in accommodating the remaining few."
    },
    "1005/1005.0985_arXiv.txt": {
        "abstract": "{% We present results of our 5-years-long program of ground-based spectroscopic and photometric observations of individual Kepler asteroseismic targets and the open clusters NGC\\,6866 and NGC\\,6811 from the Kepler field of view. We determined the effective temperature, surface gravity, metallicity, the projected rotational velocity and the radial velocity of 119 Kepler asteroseismic targets for which we acquired high-resolution spectra. For many of these stars the derived atmospheric parameters agree with $T_{\\rm eff}$, $\\log g$, and [Fe/H] from the Kepler Input Catalog (KIC) to within their error bars. Only for stars hotter than 7000\\,K we notice significant differences between the effective temperature derived from spectroscopy and $T_{\\rm eff}$ given in the KIC. For 19 stars which we observed photoelectrically, we measured the interstellar reddening and we found it to be negligible. Finally, our discovery of the $\\delta$ Sct and $\\gamma$ Dor pulsating stars in the open cluster NGC\\,6866 allowed us to discuss the frequency of the occurrence of $\\gamma$ Dor stars in the open clusters of different age and metallicity and show that there are no correlations between these parameters. } ",
        "introduction": "Our program of ground-based spectroscopic and photometric observations of stars selected for the asteroseismic targets for the Kepler space telescope by the Kepler Asteroseismic Science Consortium KASC\\footnote{Kepler Asteroseismic Science Consortium (KASC) is a group of collaborating scientists and/or institutions established to accomplish the activities of the Kepler Asteroseismic Investigation (KAI), represented by Ronald Gilliland (see http://astro.phys.au.dk/KASC).} was started in 2005 at the Osservatorio Astrofisico di Catania, OACt, (the {\\it M.G. Fracastoro\\/} station, Mt.\\ Etna, Italy) and is continued since then. Apart from the OACt, we perform spectroscopic and photometric observations of Kepler asteroseismic targets at the F.\\ L.\\ Whipple Observatory, FLWO, (Mount Hopkins, Arizona, USA), and the Astrophysical Observatory of the University of Wroc\\l{}aw in Bia\\l{}k\\'ow (Poland).  At the OACt, we use a 91-cm telescope, at FLWO, a 1.5-m telescope, and at the Bia\\l{}k\\'ow Observatory, a 60-cm telescope. We make use also of the archival data collected at the 1.5-m telescope at the Oak Ridge Observatory, ORO, (Harvard, Massachusetts, USA) and at the Multiple Mirror Telescope, MMT, before it was converted to the monolithic 6.5-m mirror.  Using the spectroscopic data acquired at OACt, FLWO, ORO, and MMT, we aim at the determination of the atmospheric parameters of the program stars, i.e., the effective temperature, $T_{\\rm eff}$, surface gravity, $\\log g$, and metallicity, $\\rm [Fe/H]$, and measuring the projected rotational velocity, $v\\sin i$, and the radial velocity, $v_r$, of the stars.  At the Bia\\l{}k\\'ow Observatory, we perform photometric time-series observations of the open clusters NGC\\,6866 and NGC\\,6811 in  the Kepler field of view, aiming at the discovery of new pulsating stars, and the determination of the degree of the modes of their pulsations.  Both sites, the Bia\\l{}k\\'ow Observatory and the OACt, took part in the international multi-site photometric campaign on NGC\\,6866 launched in 2009, and will take part in a similar campaign on NGC\\,6811 which will be launched in 2010 (for the details, see Uytterhoeven et al. 2010b.)  \\begin{figure} \\includegraphics[width=60mm,height=80mm,angle=270]{mol_fig2.ps} \\caption{The Kepler asteroseismic targets observed at the Osservatorio Astrofisico di Catania, the F.\\ L.\\ Whipple Observatory, and the Oak Ridge Observatory. Indicated are the borders of the 42 Kepler CCDs and the borders of the constellations.} \\label{targets} \\end{figure} ",
        "conclusions": ""
    },
    "1005/1005.2722_arXiv.txt": {
        "abstract": "Solar flares are currently understood as the explosive release of energy stored in the form of stressed magnetic fields.  In many cases, the released energy seems to take the form of large numbers of electrons accelerated to high energies (the nonthermal electron ``thick target'' model), or alternatively plasma heated to very high temperatures behind a rapidly moving conduction front (the ``thermal'' model).  The transport of this energy into the remaining portion of the atmosphere results in violent mass motion and strong emission across the electromagnetic spectrum.  Radiation processes play a crucial role in determining the ensuing plasma motion.  One important phenomenon observed during flares is the appearance in coronal magnetic loops of large amounts of upflowing, soft X-ray emitting plasma at temperatures of $1-2\\ \\times 10^7$ [K].  It is believed that this is due to chromospheric evaporation, the process of heating cool ($T \\sim 10^4$ [K]) chromospheric material beyond its ability to radiate.  Detailed calculations of thick target heating show that if nonthermal electrons heat the chromosphere directly, then the evaporation process can result in explosive upward motion of X-ray emitting plasma if the heating rate exceeds a threshold value.  In such a case, upflow velocities approach an upper limit of roughly $2.35 c_s$ as the heating rate is increased beyond the threshold, where $c_s$ is the sound speed in the evaporated plasma.  This is known as explosive evaporation.  If the flare heating rate is less than the threshold, evaporation takes place indirectly through thermal conduction of heat deposited in the corona by the energetic electrons.  Upflows in this case are roughly 10 to 20\\% of the upper limit.  Evaporation by thermal model heating always takes place through thermal conduction, and the computed upflow speeds seem to be about 10\\% to 20\\% of the upper limit, independent of the energy flux.  The pressure increase in the evaporated plasma for either the thick target or thermal model leads to a number of interesting phenomena in the flare chromosphere.  The sudden pressure increase initiates a downward moving ``chromospheric condensation'', an overdense region which gradually decelerates as it accretes material and propagates into the gravitationally stratified chromosphere.  Solutions to an equation of motion for this condensation shows that its motion decays after about one minute of propagation into the chromosphere.  When the front of this downflowing region is supersonic relative to the atmosphere ahead of it, a radiating shock will form.  If the downflow is rapid enough, the shock strength should be sufficient to excite UV radiation normally associated with the transition region, and furthermore, the radiating shock will be brighter than the transition region.  These results lead to a number of observationally testable relationships between the optical and ultraviolet spectra from the condensation and radiating shock. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.5102_arXiv.txt": {
        "abstract": "\\noindent  We calculate solar models including dark matter (DM) weakly-interacting massive particles (WIMPs) of mass 5-50\\gev\\;and test these models against helioseismic constraints on sound speed, convection zone depth, convection zone helium abundance, and small separations of low-degree p-modes.   Our main conclusion is that both direct detection experiments and particle accelerators may be complemented by using the Sun as a probe for WIMP DM particles in the 5-50\\gev\\;mass range. The DM most sensitive to this probe has suppressed annihilations and a large spin-dependent elastic scattering cross section. For the WIMP cross-section parameters explored here, the lightest WIMP masses $<$\\,10\\gev\\;are ruled out by constraints on core sound speed and low-degree frequency spacings.  For WIMP masses 30-50\\gev, the changes to the solar structure are confined to the inner 4\\% of the solar radius and so do not significantly affect the solar p-modes. Future helioseismology observations, most notably involving g-modes, and future solar neutrino experiments may be able to constrain the allowable DM parameter space in a mass range that is of current interest for direct detection. ",
        "introduction": "\\noindent We explore the role of DM WIMPs in modifying the thermal gradient of the Sun. A similar study involving standard WIMPS of mass $50\\gev$  or larger has been performed by Bottino et al. \\cite{Bottino}. However, here we wish to consider the effects of low mass WIMPs, with masses as low as $5\\gev$, with large trapped abundances within the Sun.  To affect the Sun's thermal gradient, we need large elastic scattering rates. The solar sound speed can be  affected, and helioseismology has been  proposed as providing a possible constraint on supersymmetric WIMPs \\cite{lopesa,lopesb, lopesc}. For the range in masses in which we are interested, the limits on the size of spin-independent WIMP elastic cross sections from CRESST \\cite{cresst}, XENON-10 \\cite{xenon10}, and XENON-100 \\cite{xenon100} are already quite stringent making it unlikely for helioseismology to provide further restrictions. Moreover, recent COUPP results \\cite{Behnke} (especially at masses $\\gsim 10$\\gev) and the results from PICASSO \\cite{picasso} similarly restrict spin-dependent interactions. However, these limits become weaker as the DM mass is decreased, especially for masses of around $5$\\gev or less.   Detailed solar models \\cite{Bottino} show that a WIMP signal is only possible for cross sections in a limited range. For a DM particle mass of $50 \\gev$ the relevant effective cross section for these signals was found to be of order $10^{-35}\\,{\\rm cm}^2$. The limits on spin-dependent scattering from direct detection experiments, e.g., COUPP \\cite{Behnke, coupp} restrict the spin-dependent elastic scattering cross section for a $50 \\gev$ DM particle to less than $\\sim10^{-37}\\,{\\rm cm}^2$.  Moving to lighter masses alleviates these limits somewhat, and for around $5 \\gev$ we are in the interesting region of around $10^{-35}\\,{\\rm cm}^2$ \\cite{picasso}. In addition, given the astrophysical uncertainties that can affect these limits (e.g.\\,\\cite{green}) it is intriguing to ask whether helioseismology can complement direct detection limits at lower masses.  Recently, there has been interest in exploring the low DM mass regime as a possible way to consistently combine the results from the DAMA direct detection experiment \\cite{dama} with those from others such as CDMS \\cite{cdms} and CoGeNT \\cite{cogent} (see e.g., \\cite{lowmasssector}). Although we do not attempt to do the same here, we simply note that this mass regime is of great interest with upper limits on the spin-dependent elastic scattering cross section for low-mass DM, reaching $\\sim10^{-32}\\,{\\rm cm}^2$ \\cite{Kopp} for certain models and assumptions.  Solar effects are most pronounced for DM with a suppressed annihilation cross section such that after capture by the Sun, the DM candidates do not annihilate quickly. A prominent example of this is asymmetric DM where annihilation is completely suppressed.  In this paper, we outline a class of models for WIMPs that are capable of modifying the temperature profile in the core of the Sun, and illustrate their effects on helioseismology and neutrino fluxes. The accumulation of these WIMPs in the solar core results in significant energy transfer to solar protons. We note that the effect is not large enough to account for discrepancies between observations and helioseismology for models that also predict the observed neutrino flux.  However, given the current debate about the appropriate element abundances to be adopted in solar models \\cite{serenelli}, this effect may still play a role and should be included in the models.  Indeed the recently revised solar abundances result in solar models that cannot reproduce currently observed helioseismic data \\cite{AGS05, AGSS09}.  The effects on the Sun of low-mass asymmetric WIMPs possessing large self-interactions have been considered in \\cite{Frandsen}. While in \\cite{Frandsen} the authors focus on WIMPs with spin-independent interactions, in this study we specifically focus on WIMPs with purely axial interactions and consequently only spin-dependent elastic scattering.  In the following sections, we outline the main features of DM models that can be potentially probed using solar properties and explore their effect on solar models. ",
        "conclusions": "Our main conclusion is that both direct detection and accelerator probes may be complemented by using the Sun as a probe of DM.  Models of DM that have large spin-dependent interactions and an intrinsic asymmetry that prevents post freeze-out annihilations can significantly lower the central temperature of the Sun as well as the resulting $^8$B neutrino flux.  For WIMP masses $m_{\\chi}>10\\gev$, the presence of WIMPs does not significantly affect currently available helioseismic constraints.  However, for WIMP masses of 10\\,GeV or lighter, constraints on sound speed and small frequency separations between $\\ell=0$ and $\\ell=2$ p-modes can begin to constrain and rule out the presence of WIMPs with the cross sections utilised here.  Our study is motivated in a large part by the recently revised solar abundances \\cite{AGS05, AGSS09} which result in solar models that cannot reproduce the currently observed helioseismic data, with numerous attempts to restore agreement being met with only partial success (see e.g., \\cite{BA08, GM10}).  This means that additional physics must be incorporated into solar modelling, and dark matter is among the options that merit detailed consideration.  Since the original submission of our paper in May 2010, an additional paper appeared on solar models including WIMPs and the implications for reconciling the new solar abundances with helioseismology \\cite{Taoso10}.  In agreement with \\cite{Taoso10}, our explorations to date do not show any realistic path in which the inclusion of WIMPs will mitigate this problem.  Even for the large interaction cross section and small annihilation cross section considered here, the inclusion of WIMPs of mass $m_{\\chi}>10\\gev$ has little effect on presently observable helioseismic signatures.  The inclusion of WIMPs with masses of 10\\,GeV or lighter worsens the agreement with the helioseismically-inferred sound speed at radii $0.1\\le r\\le0.2\\,R_{\\odot}$, only slightly deepens the predicted convection-zone depth, and introduces a trend with radial order in the low-degree p-mode small separations that is not observed in the data. Our primary new result is that WIMP masses of $\\sim 5$ GeV may be excluded for spin-dependent interactions in a specified cross-section range, thereby complementing direct detection experiments in a region that they access only with great difficulty provided the WIMPs annihilation cross section is suppressed.  While here we do not discuss whether these WIMP models could accommodate measurements of the $^8$B neutrino flux from current solar neutrino experiments such as Super-Kamiokande III \\cite{Yang:2009hp}, SNO \\cite{Collaboration:2009qz} or Borexino \\cite{Collaboration:2008mr}, with precisions of $\\sim$10\\% and theoretical expectations of up to $\\sim$20\\% depending on the solar composition \\cite{Bottino, Frandsen}, a more detailed study of the low-mass region of WIMP parameter space and its consistency with current experimental data is deferred to a later paper. There we will address the question of whether future helioseismic observations, most notably using g-modes, and solar neutrinos, may be able to constrain the allowable DM parameter space in a mass range that is of current interest for direct detection.  Finally, we note that for solar mass stars near the centre of the Galaxy, where the WIMP density is enhanced by up to some 6 orders of magnitude relative to that in the solar neighbourhood, the effect  of the redistribution of energy in the stellar core may generate a significant reduction of the main-sequence lifetimes.  We leave an investigation of this scenario to our future work."
    },
    "1005/1005.0511_arXiv.txt": {
        "abstract": "We compare the radial locations of 178 core-collapse supernovae to the $R$-band and H$\\alpha$ light distributions of their host galaxies. When the galaxies are split into `disturbed' and `undisturbed' categories, a striking difference emerges. The disturbed galaxies have a central excess of core-collapse supernovae, and this excess is almost completely dominated by supernovae of types Ib, Ic and Ib/c, whereas type II supernovae dominate in all other environments. The difference cannot easily be explained by metallicity or extinction effects, and thus we propose that this is direct evidence for a stellar initial mass function that is strongly weighted towards high mass stars, specifically in the central regions of disturbed galaxies. ",
        "introduction": "\\label{sec:intro}  Following the pioneering work of \\citet{lars78}, many studies have confirmed that tidal disturbance following galaxy interactions is an efficient trigger of star formation in galaxies \\citep[e.g.][]{jose84,kenn84}. Such star formation frequently takes the form of centrally-concentrated nuclear starbursts \\citep{jose85}, fuelled by the central concentrations of molecular gas found to occur naturally in simulations of highly-disturbed systems \\citep{barn91,miho96}. The strength of the link between starbursts and interactions was highlighted by the finding that almost all of the `ultra-luminous infrared galaxies' (ULIRGs) display signs of interactions or mergers \\citep{sand88,born99}, and by correlations between galaxy-galaxy separations and starburst strength \\citep{bart00}. Even minor mergers with low-mass companions have been shown through simulations to result in significant nuclear star formation activity \\citep{miho94}.  Several early studies of nuclear starbursts suggested that this star formation might require a top heavy initial mass function (IMF), preferentially producing high mass stars \\citep{riek80, doyo92}. There is theoretical support for this suggestion, with simulations showing that an IMF weighted to high-mass stars naturally arises in high-density regions, due to feedback processes heating the gas. In a recent study, \\citet{krum10} have demonstrated that such regions should have a high-mass stellar fraction at least 1.7 times larger, and possibly much more, than lower density, more quiescent regions.  However, the observational evidence for this variation has to date proved controversial (see \\citealt{bast10} for a recent review). Some studies have found indirect evidence for top-heavy IMFs with, for example, \\citet{riek93} concluding that the nearby starburst galaxy M82 requires an IMF biased to high mass stars to explain its emission line ratios and total luminosity. Similar techniques have been used for NGC 3256 (an ongoing merger with a `super-starburst') which have again shown indications of a modified IMF with an excess of high mass stars \\citep{doyo94}. \\citet{gibs97} showed that, in order to reproduce the observed colour-luminosity relation of elliptical galaxies, an IMF much flatter than that of \\citet{salp55} needed to be adopted. \\citet{baug05} had to employ a top heavy IMF for the starbursts powering the distant population of highly luminous submillimetre galaxies in order to explain the number counts of these systems.  Finally, \\citet{bras07} studied nine interacting galaxies from the $Chandra$ survey and found that highly disturbed systems showed a strongly enhanced infrared luminosity compared to that expected from the x-ray emission, again suggesting the need for a top-heavy IMF.  More direct evidence of a variation in IMF has been found for the resolved stellar population of the young Arches cluster in the Galactic Centre. \\citet{fige99,stol02,paum06,espi09} all find evidence for stellar mass functions weighted towards high-mass stars in this cluster or the general Galactic Centre region. Such mass functions are parametrized as an IMF that is either much flatter than that found by \\citet{salp55}, or having a higher mass turnover than is found in the function for field stars.  One possible tracer of the IMF that has not been fully exploited to date is the relative numbers of core-collapse supernovae (CCSNe) of different types. Their short progenitor lifetimes and high luminosities make them powerful indicators of recent or ongoing star formation, and indeed they provide the only direct tracer of recent star formation within unresolved stellar populations. Recent advances in the understanding of supernovae and their progenitors raise the possibility that they can provide information on the initial mass function of a young stellar population. Theoretical models of single star progenitors predict that SNII should have lower mass progenitors than SNIb or SNIc \\citep{hege03, eldr04}. This has received observational support from studies of the strength of association with H$\\alpha$ emission \\citep{ande08}, confirming that SNII have the lowest mass CC progenitors, but additionally indicating that the SNIc have still higher mass progenitors than the SNIb. The existence of this II-Ib-Ic progenitor mass sequence allows information on the IMF of the stellar population in the SN environments to be derived from the relative numbers of type II, Ib and Ic supernovae.  \\citet{petr95} studied the distribution of SNe events in 32 interacting systems containing 12 known core collapse SNe. They found that the radial distribution of these core collapse events showed a higher concentration towards the nuclear regions of the interacting galaxies when compared to isolated galaxies. This confirmed the enhanced star formation around the central regions of the systems, but the sample was too small to analyse the separate types of CCSNe.  This paper will therefore use a larger sample of local CCSNe to explore the IMF in nuclear starbursts, resulting from galaxy disturbance, by studying the ratio of type II/Ibc SNe in both disturbed and undisturbed host galaxies. Throughout this paper, we use `Ibc' to encompass all SNe with classifications of Ib, Ic or Ib/c.  The structure of the paper is as follows: In Section~\\ref{sec:sample} we will define and discuss the sample used throughout this work. Section~\\ref{sec:results} will describe the results on the radial distributions, for disturbed and undisturbed hosts and looking separately at type II and Ibc SNe. In Section~\\ref{sec:disc} we discuss the possible interpretations of our results, in terms of metallicity, extinction and IMF effects. Finally, Section~\\ref{sec:conc} contains a summary of our conclusions. ",
        "conclusions": "\\label{sec:conc}  We have analysed the spatial distribution of 178 CCSNe within a sample of host galaxies with recession velocities less than 6000~km/s. Host galaxies were classified by eye according to whether they show disturbance due to strong tidal interactions or mergers. The main results are as follows:  \\begin{itemize} \\item CCSNe of all types show a strong degree of central concentration in the disturbed galaxies, probably as a result of nuclear starbursts in these galaxies. \\item This central excess is dominated by SNIbc. \\item The SNIbc in disturbed galaxies are more centrally concentrated than the H$\\alpha$ emission. \\item The SNIbc excess cannot easily be explained in terms of metallicity effects, extinction, or central incompleteness of SNe. \\item Our preferred explanation of the SNIbc excess is that the central regions of the disturbed galaxies are dominated by nuclear starbursts with IMFs biased towards high mass stars, although metallicity, binarity and stellar rotation may also play a role. \\end{itemize}  \\ack  This paper has made use of data provided by the Central Bureau for Astronomical Telegrams. We would like to acknowledge members of staff at the Astrophysics Research Institute, in particular Sue Percival and David Bersier for their helpful comments and discussion. The authors would also like to thank the referee for a constructive and helpful report. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. The Isaac Newton Telescope is 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 Astrof\\'isica de Canarias. The Liverpool Telescope is operated on the island of La Palma by Liverpool John Moores University in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrof\\'isica de Canarias with financial support from the UK Science and Technology Facilities Council. SMH would like to acknowledge STFC for a research studentship.  \\newpage  \\begin{figure} \\noindent\\makebox[\\textwidth]{% \\includegraphics{figure1.eps}} \\caption{ Histogram showing the distribution of fractions of host galaxy $R$-band light lying within the locations of each CCSN in our undisturbed host galaxies. The top plot represents the distribution of all CCSNe, the middle SNII and the lower SNIbc.} \\label{figure1} \\end{figure}  \\newpage  \\begin{figure} \\noindent\\makebox[\\textwidth]{% \\includegraphics{figure2.eps}} \\caption{ As Figure 1, but for the disturbed host galaxies. } \\label{figure2} \\end{figure}  \\newpage  \\begin{figure} \\noindent\\makebox[\\textwidth]{% \\includegraphics{figure3.eps}} \\caption{ Histogram showing the distribution of fractions of host galaxy H$\\alpha$ light lying within the locations of each CCSN in our sample, for the undisturbed host galaxies. Again, the upper plot shows the overall CCSNe distribution, the middle SNII and the lower SNIbc. } \\label{figure3} \\end{figure}  \\newpage  \\begin{figure} \\noindent\\makebox[\\textwidth]{% \\includegraphics{figure4.eps}} \\caption{ As Figure 3, for the disturbed host galaxies. } \\label{figure4} \\end{figure}  \\newpage  \\begin{figure} \\noindent\\makebox[\\textwidth]{% \\includegraphics{figure5.eps}} \\caption{ Images of 12 of the host galaxies classified as disturbed and with centrally located CCSNe. }  \\label{figure5} \\end{figure}    \\newpage \\appendix \\setcounter{section}{1}  \\Tables \\begin{indented} \\item \\begin{longtable}{lllll} \\caption{{\\label{Table 1}}Undisturbed host galaxy sample used in this analysis. Columns represent the host galaxy, the individual SNe, the spectral classification of the SNe and the fractional $R$-band light and fractional H$\\alpha$ values for each SNe.}\\\\ \\br Host&SN&SN type&Fr($R$)&Fr(H$\\alpha$)\\\\ \\mr NGC493&1971S&IIP&0.605&0.570\\\\ NGC918&2009js&IIP&0.703&-\\\\ NGC941&2005ad&II&0.831&0.864\\\\ NGC991&1984L&Ib&0.498&0.401\\\\ NGC1035&1990E&IIP&0.272&0.363\\\\ NGC1058&1961V&II&0.968&0.931\\\\ NGC1058&1969L&IIP&1.000&1.000\\\\ NGC1058&2007gr&Ib/c&0.421&-\\\\ NGC1073&1962L&Ic&0.754&0.518\\\\ NGC1087&1995V&II&0.368&0.497\\\\ NGC1187&1982R&Ib&0.695&0.760\\\\ NGC1187&2007Y&Ib&0.981&1.000\\\\ MCG-01-09-24&2002ei&IIP&0.195&0.195\\\\ NGC1343&2008dv&Ic&0.195&0.134\\\\ UGC2906&2008im&Ib&0.682&-\\\\ UGC2971&2003ig&Ic&0.176&0.108\\\\ IC381&2001ef&Ic&0.082&0.052\\\\ NGC1832&2004gq&Ib&0.672&0.328\\\\ NGC1832&2009kr&II&0.489&-\\\\ IC2152&2004ep&II&0.461&0.560\\\\ UGC3804&2002A&IIn&0.419&0.253\\\\ NGC2551&2003hr&II&0.914&1.000\\\\ NGC2596&2003bp&Ib&0.486&0.362\\\\ UGC4436&2004ak&II&0.887&0.882\\\\ NGC2726&1995F&Ic&0.037&0.050\\\\ NGC2742&2003Z&IIP&0.675&0.736\\\\ NGC2715&1987M&Ic&0.129&0.044\\\\ UGC4904&2006jc&Ib/c&0.332&0.525\\\\ NGC2841&1972R&Ib&0.855&0.904\\\\ NGC2906&2005ip&II&0.399&0.528\\\\ UGC5249&1989C&IIP&0.017&0.058\\\\ NGC3074&1965N&IIP&0.110&0.059\\\\ NGC3074&2002cp&Ib/c&0.936&0.961\\\\ NGC3147&2006gi&Ib&0.984&0.991\\\\ NGC3184&1921B&II&0.856&0.954\\\\ NGC3184&1937F&IIP&0.808&0.930\\\\ NGC3184&1999gi&IIP&0.276&0.112\\\\ NGC3198&1966J&Ib&0.898&0.963\\\\ NGC3198&1999bw&IIn&0.745&0.755\\\\ NGC3240&2001M&Ic&0.323&0.251\\\\ NGC3294&1990H&IIP&0.156&0.125\\\\ NGC3340&2005O&Ib&0.322&0.305\\\\ \\newpage \\br Host&SN&SN type&Fr($R$)&Fr(H$\\alpha$)\\\\ \\mr NGC3340&2007fp&II&0.170&0.125\\\\ NGC3430&2004ez&II&0.788&0.833\\\\ NGC3437&2004bm&Ic&0.073&0.076\\\\ NGC3451&1997dn&II&0.872&0.946\\\\ NGC3504&2001ac&IIn&0.826&0.992\\\\ NGC3512&2001fv&IIP&0.669&0.689\\\\ NGC3556&1969B&IIP&0.197&0.494\\\\ NGC3631&1964A&II&0.915&0.992\\\\ NGC3631&1965L&IIP&0.622&0.658\\\\ NGC3631&1996bu&IIn&0.923&0.993\\\\ NGC3655&2002ji&Ib/c&0.709&0.957\\\\ NGC3683&2004C&Ic&0.532&0.545\\\\ UGC6517&2006lv&Ib/c&0.480&-\\\\ NGC3756&1975T&IIP&0.846&0.856\\\\ NGC3810&2000ew&Ic&0.261&0.147\\\\ NGC3810&1997dq&Ib/c&0.774&0.734\\\\ NGC3949&2000db&II&0.364&0.253\\\\ NGC3963&1997ei&Ic&0.197&0.053\\\\ NGC4030&2007aa&II&0.942&0.828\\\\ NGC4041&1994W&IIn&0.491&0.541\\\\ NGC4051&1983I&Ic&0.498&0.473\\\\ NGC4051&2003ie&II&0.838&0.885\\\\ IC758&1999bg&IIP&0.669&0.657\\\\ NGC4136&1941C&II&0.880&0.882\\\\ NGC4210&2002ho&Ic&0.146&0.051\\\\ NGC4242&2002bu&IIn&0.896&0.930\\\\ NGC4303&1926A&IIL&0.607&0.736\\\\ NGC4303&1961I&II&0.697&0.877\\\\ NGC4303&1964F&II&0.189&0.106\\\\ NGC4303&1999gn&IIP&0.418&0.429\\\\ NGC4303&2006ov&IIP&0.418&0.429\\\\ NGC4303&2008in&IIP&0.845&-\\\\ NGC4369&2005kl&Ic&0.271&0.540\\\\ NGC4384&2000de&Ib&0.087&0.140\\\\ NGC4451&1985G&IIP&0.138&0.212\\\\ NGC4559&1941A&IIL&0.208&0.131\\\\ UGC7848&2006bv&IIn&0.579&-\\\\ NGC4666&1965H&IIP&0.324&0.198\\\\ NGC4708&2003ef&II&0.335&0.352\\\\ NGC4725&1940B&IIP&0.675&0.802\\\\ NGC4900&1999br&IIP&0.786&0.932\\\\ NGC4961&2005az&Ic&0.426&-\\\\ NGC4981&2007C&Ib&0.320&0.236\\\\ NGC5012&1997eg&IIn&0.503&0.449\\\\ NGC5033&1950C&Ib/c&0.972&1.000\\\\ NGC5033&1985L&IIL&0.585&0.571\\\\ NGC5334&2003gm&IIn&0.480&-\\\\ NGC5371&1994Y&IIn&0.355&0.212\\\\ NGC5468&2002ed&IIP&0.811&0.791\\\\ \\newpage \\br Host&SN&SN type&Fr($R$)&Fr(H$\\alpha$)\\\\ \\mr NGC5559&2001co&Ib/c&0.618&0.497\\\\ NGC5584&1996aq&Ic&0.178&0.086\\\\ NGC5630&2005dp&II&0.534&0.590\\\\ NGC5630&2006am&IIn&0.604&0.617\\\\ NGC5673&1996cc&II&0.924&0.934\\\\ NGC5668&2004G&II&0.657&0.595\\\\ NGC5775&1996ae&IIn&0.757&0.671\\\\ NGC5806&2004dg&IIP&0.484&0.378\\\\ NGC5850&1987B&IIn&0.995&-\\\\ NGC5879&1954C&II&0.615&0.511\\\\ NGC5921&2001X&IIP&0.579&0.369\\\\ NGC6118&2004dk&Ib&0.673&0.626\\\\ NGC6207&2004A&IIP&0.729&0.660\\\\ UGC10862&2004ao&Ib&-&0.215\\\\ NGC6643&2008ij&IIP&0.519&0.620\\\\ NGC6643&2008bo&Ib&0.451&0.510\\\\ NGC6700&2002cw&Ib&-&0.642\\\\ NGC6946&2004et&II&0.975&-\\\\ NGC6951&1999el&IIn&0.320&0.259\\\\ UGC11861&1995ag&II&0.343&0.170\\\\ UGC11861&1997db&II&0.633&0.396\\\\ UGC12160&1995X&II&0.564&0.484\\\\ UGC12182&2006fp&IIn&1.000&1.000\\\\ \\br \\end{longtable} \\end{indented} \\newpage \\begin{indented} \\item \\begin{longtable}{lllll} \\caption{\\label{Table 2}Disturbed host galaxy sample used in this analysis. Columns represent the host galaxy, the individual SNe, the spectral classification of the SNe and the fractional $R$-band light and fractional H$\\alpha$ values for each SNe, as in table 1.}\\\\ \\br Host&SN&SN type&Fr($R$)&Fr(H$\\alpha$)\\\\ \\mr NGC895&2003id&Ic&0.524&-\\\\ UGC2984&2002jz&Ic&0.091&0.099\\\\ NGC1614&1996D&Ic&0.275&-\\\\ NGC1637&1999em&IIP&0.276&0.268\\\\ IC391&2001B&Ib&0.062&0.060\\\\ NGC1961&2001is&Ib&-&0.749\\\\ NGC2207&1999ec&Ib&-&0.521\\\\ NGC2207&2003H&Ib&-&0.259\\\\ NGC2146&2005V&Ib/c&0.033&0.091\\\\ ESO492-G2&2005lr&Ic&-&0.005\\\\ UGC3829&2001ej&Ib&0.152&0.391\\\\ NGC2276&1968V&II&0.699&0.790\\\\ NGC2276&2005dl&II&0.247&0.099\\\\ NGC2276&1993X&II&0.899&0.619\\\\ NGC2532&1999gb&IIn&0.485&0.443\\\\ NGC2532&2002hn&Ic&0.023&0.011\\\\ NGC2604&2002ce&II&0.381&0.560\\\\ NGC2782&1994ak&IIn&0.725&0.977\\\\ NGC2993&2003ao&IIP&0.456&0.784\\\\ NGC3169&1984E&IIL&0.684&0.731\\\\ NGC3310&1991N&Ic&0.268&0.277\\\\ NGC3323&2004bs&Ib&0.191&0.119\\\\ NGC3323&2005kk&II&0.766&0.875\\\\ NGC3367&1992C&II&0.689&0.687\\\\ NGC3367&2007am&II&0.302&0.314\\\\ NGC3627&1973R&IIP&0.471&0.566\\\\ NGC3627&1997bs&IIn&0.362&0.348\\\\ NGC3627&2009hd&II&0.496&-\\\\ NGC3690&1993G&IIL&0.464&0.744\\\\ NGC3690&1998T&Ib&0.056&0.056\\\\ NGC3690&1999D&II&0.560&0.849\\\\ NGC3786&1999bu&Ic&0.180&0.522\\\\ NGC3811&1971K&IIP&0.809&0.900\\\\ IC2973&1991A&Ic&0.742&0.588\\\\ NGC4038&2004gt&Ib/c&0.834&0.991\\\\ NGC4088&1991G&IIP&0.466&0.453\\\\ NGC4088&2009dd&II&0.100&-\\\\ NGC4141&2008X&IIP&0.194&0.085\\\\ NGC4141&2009E&IIP&0.594&0.491\\\\ NGC4254&1967H&II&0.664&0.648\\\\ NGC4254&1972Q&IIP&0.811&0.791\\\\ NGC4254&1986I&IIP&0.334&0.318\\\\ NGC4273&2008N&IIP&0.300&0.316\\\\ NGC4273&1936A&IIP&0.569&0.598\\\\ NGC4490&1982F&IIP&0.277&0.202\\\\ \\newpage  \\br Host&SN&SN type&Fr($R$)&Fr(H$\\alpha$)\\\\ \\mr NGC4568&1990B&Ic&0.302&-\\\\ NGC4568&2004cc&Ic&0.158&-\\\\ NGC4618&1985F&Ib&0.121&0.087\\\\ NGC4615&1987F&IIn&0.489&0.333\\\\ NGC4688&1966B&IIL&0.571&0.454\\\\ NGC4691&1997X&Ic&0.171&0.472\\\\ NGC5000&2003el&Ic&0.482&0.476\\\\ NGC5021&1996ak&II&0.619&0.659\\\\ MCG-04-32-07&2003am&II&0.211&0.339\\\\ NGC5194&1994I&Ic&-&0.122\\\\ NGC5194&2005cs&IIP&-&0.222\\\\ NGC5395&2000cr&Ic&0.538&0.549\\\\ NGC5480&1988L&Ib&0.230&0.369\\\\ NGC5682&2005ci&II&0.204&0.191\\\\ NGC7479&1990U&Ic&0.603&0.488\\\\ NGC7479&2009jf&Ib&0.764&-\\\\ NGC7537&2002gd&II&0.759&0.685\\\\ NGC7714&2007fo&Ib&0.377&-\\\\ UGC12846&2007od&IIP&0.945&1.000\\\\  \\br \\end{longtable} \\end{indented}                           \\clearpage \\def\\newblock{\\hskip .11em plus .33em minus .07em}"
    },
    "1005/1005.3332_arXiv.txt": {
        "abstract": "s{We present an analytical derivation of the Sachs Wolfe effect sourced by a primordial magnetic field, generated by a causal process, such as a first order phase transition in the early universe. As for the topological defects case, we apply the general relativistic junction conditions to match the perturbation variables before and after the phase transition, in such a way that the total energy momentum tensor is conserved across the transition. We find that the relevant contribution to the magnetic Sachs Wolfe effect comes from the metric perturbations at next-to-leading order in the large scale limit. The leading order term is strongly suppressed due to the presence of free-streaming neutrinos. We derive the neutrino compensation effect and confirm that the magnetic Sachs Wolfe spectrum from a causal magnetic field behaves as $\\ell(\\ell+1)\\,C^\\B_\\ell \\propto \\ell^2$ as found in the latest numerical analyses.} ",
        "introduction": "The origin of the large scale magnetic fields observed in galaxies and clusters is still unknown: one of the possible explanations is that they have been generated in the primordial universe. A primordial magnetic field of the order of the nanoGauss could leave a detectable imprint in the cosmic microwave background (CMB) anisotropies~\\cite{Kahniashvili:2006hy,Yamazaki:2008gr,PFP,Shaw:2009nf}. Here we concentrate on its effect on the temperature CMB spectrum at large scales and more particularly on the Sachs Wolfe effect. The motivation is that conflicting results exist in the literature regarding the $\\ell$-dependence of the Sachs Wolfe effect induced by a causal primordial magnetic field: the analytical analysis of \\cite{Kahniashvili:2006hy} found $\\ell(\\ell+1)\\,C^\\B_\\ell$ scaling as $ \\ell^{-1}$ or more negative, and the same result was found in the numerical calculation of \\cite{Yamazaki:2008gr}; on the other hand, \\cite{PFP,Shaw:2009nf,Caprini:2009vk} found $\\ell(\\ell+1)\\,C^\\B_\\ell$ scaling as $\\ell^2$. The aim of this paper is to explain this discrepancy analytically (for a more detailed derivation see~\\cite{cmbmag}).  We assume that a magnetic field is generated in the early universe by a sudden phase transition, as for example the electroweak (EW) phase transition. We consider a stochastic magnetic field with no background component and we suppose that the magnetic energy momentum tensor is first order in perturbation theory. We study the effect of the magnetic field on the metric and fluid (matter plus radiation) perturbations by solving analytically Einstein's and conservation equations in the long wavelength limit. We take into account the neutrinos in our derivation. In order to connect the solutions before and after the magnetic field generation, we match the geometry and the fluid variables at the phase transition time, so that the induced three metric and the extrinsic curvature are continuous~\\cite{Deruelle:1997py}. This implies the conservation of the total energy momentum tensor across the phase transition, and it completely determines the metric and fluid perturbation variables after the magnetic field generation. Before neutrino decoupling, we find that at leading order in the large scale expansion $k/\\mathcal{H}\\ll 1$, the metric perturbation $\\Phi$ is proportional to $\\Phi \\propto \\Pi_\\B (\\HH/k)^2$, where $\\Pi_\\B$ is the magnetic field anisotropic stress. This induces a contribution in the CMB spectrum scaling as $\\ell(\\ell+1)\\,C^\\B_\\ell\\propto  \\ell^{-1}$, and consistent with~\\cite{Kahniashvili:2006hy,Yamazaki:2008gr}. However, once neutrinos decouple and start free-streaming, they acquire a non-zero anisotropic stress, which acts to compensate and reduce the magnetic field one~\\cite{Shaw:2009nf,Kojima:2009gw}. We demonstrate that this compensation drastically reduces the leading order contribution to the CMB spectrum, and that the dominant contribution becomes the one from the next-to-leading order in the $k/\\mathcal{H}\\ll 1$ expansion, which induces then $\\ell(\\ell+1)\\,C^\\B_\\ell \\propto \\ell^2$, as found in~\\cite{PFP,Shaw:2009nf,Caprini:2009vk}. ",
        "conclusions": "In this work we present an analytical computation of the Sachs Wolfe effect induced by a primordial magnetic field. We have restricted our analysis to a magnetic field generated by a causal process, such as a first order phase transition. In this case, the initial conditions for the metric and fluid variables are determined unambiguously by imposing conservation of the total energy momentum tensor across the transition. Using these initial conditions, we have computed analytically the leading order and next-to-leading order magnetic contribution to the Sachs Wolfe effect. We have found that the leading order contribution is sourced only by the magnetic field anisotropic stress, and leads to a CMB spectrum scaling as $1/\\ell$. However, this contribution is strongly suppressed once the magnetic field anisotropic stress is compensated by the one of the neutrinos. As a consequence, the dominant contribution to the Sachs Wolfe is the next-to-leading order one, that generates a CMB spectrum scaling as $\\ell^2$. Our analytical work solves therefore the discrepancy regarding the $\\ell$-dependence of the magnetic Sachs Wolfe."
    },
    "1005/1005.1876_arXiv.txt": {
        "abstract": "Long-duration gamma-ray bursts(LGRBs) are believed to be linked with the star formation. We adopt a galactic evolution model, in which the star formation process inside the virialized dark halo at given redshift can be achieved. In this paper, the gamma-ray burst(GRB) host galaxies are assumed to be the star-forming galaxies within the small dark halos. The star formation rates(SFRs) in the host galaxies of LGRBs at different redshifts have been derived from our model with the galactic evolutionary time about a few times of $10^7$ yr and the dark halo mass of about $5\\times 10^{11}M_\\odot$. The related stellar masses, luminosities and metallicities of these hosts are estimated as well. We further calculate the X-ray and optical absorption of GRB afterglow emission. From our model calculation, at higher redshift, the SFR of host galaxy is larger, the absorption in X-ray band and optical band of GRB afterglow is stronger, in the condition that the dust and metal components are released locally, surrounding the GRB environment. These model predictions are compared with the {\\it Swift} and other observational data. At lower redshift $z<1$, as the merger and interaction of some host galaxies are involved, one monolithic physical process is not sufficient to fully explain all kinds of observed phenomena. ",
        "introduction": "Gamma-ray burst(GRB), the most violent explosion cosmic source, has been identified as the cosmological event since 1997(van Paradijs et al. 1997; Metzger et al. 1997). Recently GRB 090423 has been explored at high redshift above 8(Salvaterra et al. 2009a; Tanvir et al. 2009). The long-duration GRB(LGRB) progenitors are proposed to be the massive collapsing stars (Woosley 1993; Kumar, Narayan \\& Johnson 2008). Some long bursts have been observed to be associated with supernova events(Hjorth et al. 2003; Stanek et al. 2003; Malesani et al. 2004; Mazzali et al. 2006; Xu et al. 2008), hence having a common star-forming origin(Paczynski 1998). Indeed, long GRBs can be found in the star formation galaxies and these galaxies are dominated by the young stellar population(Christensen et al. 2004). In general, GRBs favor a metal-poor environment(Fynbo et al. 2006; Kewley et al. 2007) and the hosts have low stellar masses(Wiersema et al. 2007). Jakobsson et al. (2005) proposed that GRB host galaxies, at least those high redshift($z>2$) hosts, trace the star formation of the universe in an unbiased way. The high global star formation rate(SFR) history at redshift larger than 6 (Hopkins \\& Beacom 2006; Yan et al. 2009) indicates the possibility of high-redshift GRB production and the detection of host galaxies. From the research of Y\\\"{u}ksel et al. (2008) and  Kistler et al. (2009), there could be a link from star formation to the GRB production in the high redshift universe, in which the GRB luminosity function is involved. Moreover, the evolution of the GRB luminosity function has been investigated by Salvaterra et al. (2009b). All of these evidence provide the strong clue to study the intrinsic link from SFR to GRB production and the possible evolutionary properties of GRBs and their hosts.  The grains and metals produced by the host galaxy will take effects on the GRB afterglow emissions. Thus, the GRB progenitors and their environments can be expressed by the absorption features of GRB afterglows. The heavy attenuation in the X-ray band has been given in the statistic results from Campana et al. (2010), indicating a dense surrounding environment of those GRBs. In the mean while, it is also interesting to understand whether this kind of strong attenuation is intrinsically evolved with redshift. On the other hand, the characteristics of the corresponding absorption in the optical band are still under debate. Although the approximate dust extinction law of GRB host galaxies has been given by Chen, Li \\& Wei (2006) and Li et al. (2008), in order to have an explanation of dust obscuration and especially to interpret some X-ray detected but optical faint bursts(so-called dark bursts, Akerlof \\& Swan 2007; Kann et al. 2007; Perley et al. 2009), the physical origin associated with the star formation and galactic evolution should be studied in an unified scenario.  In this paper, we specify one physical model of star-forming and metal-poor galaxies being as the hosts of long GRBs, exploiting the physical recipes from Granato et al. (2004). In the general scenario of Granato et al. (2004), at each redshift bin, the SFR and galaxy mass in the given dark halo potential well have been calculated, with the effects on the kinetic feedback of supernova and central black hole. Under this framework, the different evolutionary stages of galaxies and the central black holes with different physical conditions have been investigated(e.g., Cirasuolo et al. 2005 about the properties of E/S0 galaxies; Lapi et al. 2006 about the active galactic nucleus luminosity function; Granato et al. 2006 about the submillimeter galaxies). In particular, Mao et al. (2007) calculated the UV luminosities and the relative dust attenuation in the star-forming and metal-poor galaxies, Lapi et al (2008) estimated the long GRB progenitor rates and redshift distribution. Since the updated X-ray/optical observations on the GRB afterglows and host galaxies have been performed sequentially by Castro Cer\\'{o}n et al. (2008), Evans et al. (2009), Savaglio, Glazebrook \\& Le Borgne (2009), Levesque et al. (2009a) and Fynbo et al. (2009), in this context, it is necessary to further compare some properties calculated by our model with these updated observational data. We extend the former calculation from Mao et al. (2007), attempting to understand the physical origin of the long GRB production and the GRB environment, especially, we reveal that some properties from afterglow emissions and GRB hosts have shown the possible intrinsic cosmological evolution.  Throughout the paper, we adopt cosmological parameters: $h=0.7$, $\\Omega_M=0.3$, and $\\Omega_\\Lambda =0.7$. ",
        "conclusions": "Under the framework of galaxy formation scenario, Lapi et al. (2008) predicted the GRB progenitor rate and redshift distribution. In this paper, without the information of GRB rates and the cosmological star formation density, we attempt to reveal some properties of GRBs and their host galaxies, which have intrinsic redshift distributions. The distributions of these properties with redshift are found to be originally from the star formation in the star-forming galaxies. Given a proper galactic evolutionary time and a reasonable dark halo mass, the final results can be obtained by the model calculation. These results are compared with all kinds of observational data.  At high redshift, the GRB host galaxy has a plenty of neutral gas, suffering much violent star formation. After the short-time stellar evolution phase, the metal and dust are released by massive stars; thus, the optical and X-ray GRB emissions will have a strong attenuation locally at high redshift. The star formation activity, evolving from relatively massive hosts at high redshift to dwarf galaxies at low redshift, is a coincidence with the so-called downsizing scenario(e.g., Heavens et al. 2004). However, at lower redshift, the situation turns to be more complicated. From the morphological statistics by Conselice et al. (2005) and Wainwright et al. (2007), the GRB hosts present a broad diversity of galaxy types. About 1/3 host galaxies in the sample of Savaglio et al. (2009) are mergers, while in our model the merging and interaction processes are not taken into account. In fact, Conselice et al. (2005) found that the GRB hosts at $z>1$ are different from those at $z<1$ in terms of light concentration and the morphological size. Through the study of galaxy mass distribution, GRB hosts tracing star formation might be biased at low redshift(Kocevski et al. 2009). It is also complex that the hosts at $z<1$ are not representative of the general galaxy population(Levesque et al. 2009a). Thus, the properties of these low-redshift GRB hosts presented in this paper could not be reproduced by any monolithic process. At least, some low-redshift galaxies may undergo multiple star-forming processes during their whole lifetimes. GRB production can be accompanied with any single starburst event.  From the analysis in this paper, we see that the absorption of GRB X-ray and optical emissions is relatively strong. The strong intrinsic attenuation of GRB host galaxies may produce some dark bursts, defined by the index $\\beta_{ox}<0.5$, where $\\beta_{ox}$ is the flux density ratio between optical and X-ray bands(Jakobsson et al. 2004). Rol et al. (2005) proposed several extinction origins from their preliminary results. From our calculations, we see that the heavy attenuation may occur due to the following three possibilities: (1) the local environment of the host is metal-enriched, metallicity is higher, and/or, the host galaxy in the massive dark halo larger than $10^{12}M_\\odot$ may have strong absorption. For example, at redshift 2.5, $Z=1.0Z_\\odot$, halo mass $M_{halo}=5.0\\times 10^{12}M_{\\odot}$, after the galactic evolving time $1.0\\times 10^8~yr$, we have dust extinction $A_v=1.0$ and the corresponding X-ray absorption $N_{H,x}=4.7\\times 10^{22}cm^{-2}$; (2) the dust and metals surrounding the GRB in the host galaxy are distributed in an inhomogeneous way; there could be heavy absorption through the line of sight, but in other directions the absorption is slight. Also, in our model, we assume that the $A_v$ and $N_{H,x}$ are measured locally and do not change significantly if the dust and gas extend out to a few tens to hundreds of pc from the burst(Perna \\& Lazzati 2002, D'Elia et al. 2009); however, suppose the observed optical extinction is due to the grain absorption far beyond this local region of GRB, the $A_v$-$N_{H,x}$ correlation obtained by Schady et al. (2007, 2010) may be invalid and our calculations are strongly biased; (3) as mentioned in Section 2.2.2, the dust produced by the AGB population at high redshift should be taken into account.  In order to further understand the metal production of GRB environment, we roughly re-estimate the metallicity of GRB hosts under our framework. The mass of metal $M_{metal}=SFR\\cdot f\\cdot f_{dep}\\cdot M_{dust}/M_{star}$, where f is the ratio of massive stars to all stars, and $f_{dep}$ is the ratio of mental converted from the dust. $f=0.47$ is the case for the stars with the mass larger than $2M_\\odot$ by our adopted IMF, $f_{dep}=1.0$ means that all the dust can be transferred to metals. Metallicity is defined by $Z=M_{metal}/M_{gas}$. From the SFR calculated by Granato et al. (2004) and Mao et al. (2007), as an example, at redshift 6, we obtain the metallicity as $Z\\sim 2.75\\times 10^{-2}(M_{dust}/M_{star})$. If we take a supernova with the dust production of $10^{-3}$ solar mass(Pozzo et al. 2004), we obtain the upper limit of metallicity $Z\\sim 10^{-3}Z_{\\odot}$, which is lower than the measurement($Z>0.02Z_\\odot$) of GRB 050904(Campana et al. 2007). If we take the dust mass 0.08-0.3 solar mass per primordial massive supernova(Todini \\& Ferrara 2001), we have the result which is consistent with the observation. The estimation values of Population I/II metallicity are lower than the observational values at high redshift, meaning that the imprints from some primordial objects(Kawai et al. 2006), such as pop III stars and mini-quasars, have to be included in the possible cosmic evolution properties of these GRB host galaxies. According to this estimation, the metal-enriched environment of GRB host galaxy naturally gives the reason of the strong attenuation in X-ray and optical band measurements. In our model, the initial galactic evolutionary time of about $10^7$ yr of host galaxies is given, but the corresponding metallicity about $Z\\sim 0.3Z_\\odot$(Lapi et al. 2008) is not a necessary condition, as mentioned by Levesque et al. (2009b) that low metallicity may not be required for a relativistic explosion. With our model, the massive dark halo above $10^{12}M_\\odot$ can host the GRB galaxy in which the metallicity is relatively high, although most GRB host galaxies are inside the dark halos with the masses less than $10^{12}M_\\odot$. On the other hand, a host galaxy with a top-heavy IMF, meaning that much more massive stars are involved, can produce more metals in relatively short time during the galactic evolution phase. For instance, the Wolf-Rayet star with a mass of 80$M_\\odot$ and initial metallicity $Z=0.001$ has the possibility to self-enrich the HII region(Kr\\\"{o}ger, Hensler \\& Freyer 2006) and to produce the GRB event(Eldridge et al. 2006).  In this paper, we have calculated the SFR, galactic mass, and metallicity of the GRB host galaxies. The absorption variations with redshift in the X-ray and optical bands are presented as well. Some selection effects have been taken into account through our calculation. Other observational biases should also be considered. All the redshift measurements come from the optical observations so that some optical-faint GRBs and host galaxies are ignored. Moreover, at high redshift only most luminous galaxies with high SFRs can be detected, indicating that some low-luminosity cases are not included. However, our calculations come from the intrinsic star formation of GRB host galaxies. Thus, due to all these selection effects and observational bias mentioned in the paper, the intrinsic properties of GRB afterglows and the hosts by the model calculations have some differences to those from observations. As SFR evolution plays a dominant role in the calculations, compared to the situation at low redshift, in general, star formation in the metal-poor environment at high redshift may provide more powerful GRB explosion. Therefore, although the effective threshold is given by Kistler et al. (2009), we speculate that improving the sensitivity of detectors on the high-energy telescopes is not strongly useful to catch more high-redshift but faint GRBs, since low-energy-released GRBs are almost absent in the high-redshift universe."
    },
    "1005/1005.3042_arXiv.txt": {
        "abstract": "{In recent years mid- and far infrared spectra of planetary nebulae have been analysed and lead to more accurate abundances. It may be expected that these better abundances lead to a better understanding of the evolution of these objects.} {The observed abundances in planetary nebulae are compared to those predicted by the models of Karakas (2003) in order to predict the progenitor masses of the various PNe used. The morphology of the PNe is included in the comparison. Since the central stars play an important role in the evolution, it is expected that this comparison will yield additional information about them.} {First the nitrogen/oxygen ratio is discussed with relation to the helium/hydrogen ratio. The progenitor mass for each PNe can be found by a comparison with the models of Karakas. Then the present luminosity of the central stars is determined in two ways: first by computing the central star effective temperature and radius, and second by computing the nebular luminosity from the hydrogen and helium lines. This luminosity is also a function of the initial mass so that these two values of initial mass can be compared.} {Six of the seven bipolar nebulae can be identified as descendants of high mass stars (4M\\smallsun - 6M\\smallsun) while the seventh is ambiguous. Most of the elliptical PNe have central stars which descend from low initial mass stars, although there are a few caveats which are discussed. There is no observational evidence for a higher mass for central stars which have a high carbon/oxygen ratio. The evidence provided by the abundance comparison with the models of Karakas is consistent with the HR diagram to which it is compared. In the course of this  discussion it is shown how `optically thin' nebulae can be separated from those which are 'optically thick'.} {} ",
        "introduction": "Planetary nebulae (hereafter PNe) are an advanced stage of stellar evolution of low and intermediate mass stars. After the asymptotic giant branch (AGB) phase is completed, these stars evolve through the PN stage before ending their lives as white dwarfs. The gaseous nebula seen now as PN is the remnant of the deep convective envelope which once surrounded the core. This core is now seen as the central star of the PN. The present abundances in the nebula reveal information about the chemical processes that took place during the AGB. These processes, which have first been discussed by Iben \\& Renzini \\cite{iben} and Renzini \\& Voli \\cite{renzini}, change the abundances according to the mass of the star involved and the initial abundances in the star. Thus by investigating the PN abundances it may be possible to assign an initial mass to the star.  Models have been made of the evolution of stars of different masses. These were initiated with the discussion of Paczynski \\cite{pacz} followed by the detailed calculations of Sch\\\"{o}nberner \\cite{schon}, Vassiliadis and Wood \\cite{wood} and Bl\\\"{o}cker \\cite{blocker}. These models refer mostly to post AGB evolution. Models referring to evolution on the AGB have been made by several authors, e.g. Marigo et al. \\cite{marigo} and Karakas \\cite{karakas}. The latter models predict changes in the chemical composition which have occurred during the evolution and which have been brought to the surface and subsequently expelled as the nebula. It is these models which will be used to compare with observed PN abundances because not only do they follow a star of a given mass over its entire life, but the same is done for an entire sequence of possible masses for several different initial abundances.  The purpose of the present paper is to compare these models with the abundances which we have observed. These abundances have been determined with the help of mid and far infrared observations either from {\\em ISO} or {\\em Spitzer} and are quite accurate because they are less affected by possible temperature variations or gradients in the nebula. These observations have already been used (Pottasch \\& Bernard-Salas\\,\\cite{pobs} to better determine PN abundance gradients in the galaxy.  In an ideal case it might be expected that a comparison of models with observations will lead to: 1) knowledge of the individual properties of the central stars, and 2) confirmation or suggestion for improvement of the models.  In practice these goals are rather difficult to reach because of shortcomings of both the observations as well as the models. On the observational side are uncertainties in the effective temperature of the central star, their distances, as well as the accuracy of the measurements. The models presently available are uncertain because the physical conditions in the actual star-nebula system is poorly known. For example, the mass loss along the AGB (and post AGB) is physically not well understood and the initial conditions may not be realistic. Thus models used for comparison are taken from different authors who may use different mass loss rates. Therefore core masses are used where possible although initial masses are given for the Karakas models because the author identifies them as such.  The abundances observed are listed in Table 1. No indication is given there of the spectrum of the central star. A few of these stars are Wolf-Rayet stars for which it may be that some of the evolutionary calculations may not apply. These are the central stars of BD+30 3639, NGC\\,40, NGC\\,5315 and NGC\\,6369. This can be kept in mind when making the comparisons.  The objects were selected to be IR bright (in the diaphragm of the instrument used). This was first done with the ISO spectrometer where almost all of the usable PN spectra were investigated. Later the Spitzer IR spectra of PNe have been investigated. Most of these spectra are as bright or nearly as bright as the ISO PNe. This may at first suggest a bias toward PNe with massive central stars because these initially evolve at the highest luminosity. But the period of high luminosity is expected to be very short so that very few, if any, high mass central star PNe are expected. We therefore may expect that many low mass central star PNe have been observed, not only because of much longer evolution time but also because of the much greater number of low mass objects present. It is expected that most of the observed PNe are reasonably local objects, within a few kpc of the sun.  Nevertheless a confrontation of the models with the observations, even with these limitations, can give interesting insights into the evolution of the PN system. In Sect.\\,2 the morphology of the nebulae will be discussed, first in relation to the nitrogen/oxygen ratio observed in the nebula, and then the helium abundance will be introduced into the discussion. In Sect.\\,3 the effective temperature of the central star and the various ways of determining it will be discussed. Then the luminosity of the central stars will be discussed. Because the luminosity is dependent on the distance of the nebula this will also be discussed in this section. In addition a digression will be made into the long-standing question of whether or not a nebula is 'thick' to photons which ionize hydrogen. This can be done because two different methods of obtaining the nebular luminosity are available, one which makes no assumption concerning the nebular 'thickness' and one which is dependent on this assumption. In Sect.\\,4 we investigate whether the expected relation between abundance and luminosity can be seen. Conclusions and discussion are given in the last section. ",
        "conclusions": "We have determined the masses of a selection of 33 well-studied nebulae PNe by comparing the observed nebular abundances with that predicted by the evolutionary models of Karakas\\,\\cite{karakas}, which is the first systematic study of the evolution of lower mass stars. A secondary purpose is to see whether the masses determined in this way are consistent with the evolutionary tracks computed by Sch\\\"{o}nberner\\,\\cite{schon}, Bl\\\"{o}cker\\,\\cite{blocker} and Vassiliadis \\& Wood\\,\\cite{wood}. The abundances used in this comparison are helium, oxygen, nitrogen and carbon. The morphology is also considered in this comparison but in a simplified form. The PNe are divided into two categories: bipolar and elliptical (including round nebulae). The result of the comparison of He vs N/O (section 2.2) indicate four PNe have core masses greater than 0.9M\\smallsun; three of these nebulae are bipolar (NGC6302, NGC6537 and He2-111). There are also four nebulae which appear to have a slightly lower initial mass, between 0.85M\\smallsun~and 0.9M\\smallsun~and have also undergone hot-bottom burning.  Again three of these PNe are bipolar (NGC2440, Hb5 and Hu1-2). Thus 6 of the 7 bipolar nebular in our selection have abundances which indicate that they originated from initially high mass stars. The only bipolar nebula which appears to be an exception is NGC6445 which has a high He/H ratio but a low N/O ratio which is difficult to understand on the basis of the models of Karakas.  NGC5315 appears to be an elliptical PN with a high mass. There are three other elliptical nebulae which on the basis of a rather high N/O ratio seem to originate from stars of rather high mass. These are NGC6153, NGC6886 and NGC2392. A unlikely explanation is that these PNe are really bipolar seen edge on making them look round. Two of these PNe do not have a high He/H ratio however, making this interpretation appear too simplified. All the other elliptical PNe have N/O and He/H ratios which do not substantially differ from solar and therefore have a core mass of less than approximately 0.7M\\smallsun, using the evolutionary models of Karakas\\,\\cite{karakas}. We cannot determine the mass more precisely for these nebulae, except to say that in the models of Karakas those PNe with a high carbon abundance (C/O$\\geq$1) have the higher mass.  An HR diagram for these nebulae was then constructed in order to see if the position of the PNe on this diagram can confirm the conclusions drawn from the comparison of the abundances with the predictions of Karakas. The effective temperature of the central stars is determined, usually with an error that is less than 10\\%. The determination of the luminosity can be made in two ways. First by computing it from the temperature and radius of the central star, and secondly by finding it from the nebular emission in both the hydrogen and helium lines. The second method has the additional assumption that the nebula is 'optically deep' to ionizing radiation of hydrogen or helium. We find that the luminosity determined from the ionized helium line is usually approximately equal to the luminosity found from the central star, indicating that the nebula absorbs all radiation which can doubly ionize helium. For about 30\\% of the nebulae the hydrogen ionizing radiation gives a lower luminosity indicating that in these cases the nebula is 'optically thin' to this radiation. In this way a method has been found to determine which nebulae are `optically deep' to the various radiation fields. Probably the largest uncertainty in the luminosity is the distance determination so that a subsection is devoted to a discussion of the distances used. The resultant HR diagram is shown in Fig.4. The high masses found from the abundances for the six bipolar PNe are consistent with their position on the HR diagram. The position of the seventh bipolar PN, NGC6445, is consistent with it being a high mass object, leaving its rather low N/O ratio as a problem. The position of NGC6886 is consistent with a rather high mass. Most of the elliptical PNe have positions consistent with low core masses. The biggest problems are: 1)  NGC7027, whose HR position indicates a high mass while its abundances give a low  mass; 2) NGC5315 whose position indicates a low mass while its abundance gives a high mass. M1-42 is also a problem but its distance is very poorly known and some of the observations are not very good. NGC2392 has a different problem since the luminosity determined from the star is improbably high. The effective temperature that we have determined for this central star does not apply to the star measured in the visual. The most likely solution is that this central star is a binary and the secondary, which is unseen in the visual, is of a much higher temperature and is responsible for the high degree of ionization found. The possible presence of a hotter star has been suspected earlier; Ciardullo et al.\\cite{ciar} have looked for it but have found only a very faint nearby star. If this is the source of ionization it cannot be a main sequence star because if this was so it would be too distant to be associated with the nebula. But the ionization source could be a star which is too close to the bright star to be observed.  In general it appears that the initial masses as determined from the observed abundances in conjunction with the models of Karakas\\,\\cite{karakas} are consistent with the initial masses predicted using the evolutionary models of Sch\\\"{o}nberner\\,\\cite{schon}, Bl\\\"{o}cker\\,\\cite{blocker} and Vassiliadis \\& Wood\\,\\cite{wood}. But being consistent is only a first step and both the models and the observed abundances should both be improved. Furthermore the various cases which appear to give inconsistent results must be understood before we can speak of agreement."
    },
    "1005/1005.3797_arXiv.txt": {
        "abstract": "Direct detection of dark matter (DM) requires an interaction of dark matter particles with nucleons. The same interaction can lead to dark matter pair production at a hadron collider, and with the addition of initial state radiation this may lead to mono-jet signals. Mono-jet searches at the Tevatron can thus place limits on DM direct detection rates.  We study these bounds both in the case where there is a contact interaction between DM and the standard model and where there is a mediator kinematically accessible at the Tevatron.  We find that in many cases the Tevatron provides the current best limit, particularly for light dark matter, below $\\sim 5$ GeV, a and for spin dependent interactions. Non-standard dark matter candidates are also constrained. The introduction of a light mediator significantly weakens the collider bound. A direct detection discovery that is in apparent conflict with mono-jet limits will thus point to a new light state coupling the standard model to the dark sector. Mono-jet searches with more luminosity and including the spectrum shape in the analysis can improve the constraints on DM-nucleon scattering cross section. ",
        "introduction": "\\label{sec:intro}  From astronomical and cosmological observations it is now clear that $\\sim 25\\%$ of the matter-energy content of the universe if made up by dark matter (DM).  Although DM has so far only been observed through its gravitational interactions the quest for a more direct observation of DM is taking place simultaneously on many fronts.  Indirect searches look for signals of standard model (SM) particle production from DM annihilations in our galaxy, direct searches look for interactions of DM with SM particles in underground detectors and colliders attempt to produce the DM and measure it.  We will concentrate here on direct detection and collider searches.   If dark matter is to be observed in direct detection searches it must couple to quarks or gluons~\\footnote{DAMA is an exception as, unlike other experiments, it does not distinguish between nuclear and electron recoils of DM.}. The same couplings lead to direct DM production at hadronic colliders such as the Tevatron, and we wish to investigate the connection between the two types of search.  We will do so in a model independent fashion~\\cite{Birkedal:2004xn}; we will assume that the DM is fermionic and that there is some massive state whose exchange couples DM to quarks. The mediator may be a SM gauge boson, the Higgs or a new particle (if the new particle is very heavy we can describe its effects with an effective contact operator). Although the processes that give direct detection and those that give DM production occur through s- and t-channel exchange  of the same mediator, the regimes probed in the two types of experiment are very different.  The momentum exchange during a DM-nucleus recoil is $\\sim 100\\ \\mev$ whereas at the Tevatron the typical momentum exchange is $10-100\\ \\gev$.  This leads to two interesting regimes to consider when comparing bounds from the two types of experiments: heavy mediators $M\\gtrsim  100\\ \\gev$ and light mediators $M\\lesssim 100\\ \\gev$.   The momentum exchange at direct detection experiments is sufficiently low that for all but the lightest mediators below ${\\cal O}(100~\\mbox{MeV})$, which we do not consider here, the mediator can effectively be integrated out and the scattering rate in both regimes scales as, \\be \\sigma_{\\rm DD}\\sim g_\\chi^2\\, g_q^2\\, \\frac{\\mu^2}{M^4}\\,, \\ee where, for simplicity, we have ignored form factors and possible momentum and velocity dependence in the cross section. Here, $g_\\chi$ and $g_q$ are couplings of the mediator to DM and quarks. $\\mu$ is the reduced mass of the DM-nucleon system.  In contrast the two regimes behave very differently at colliders.  Concentrating on direct production of a pair of DM particles and an initial state emission of a jet, we estimate the mono-jet + $\\slsh{E_T}$ partonic production cross section in the two cases to be \\be \\sigma_{1j} \\sim \\begin{cases} \\alpha_s \\,g_\\chi^2\\, g_q^2\\,\\frac{1}{p_T^2} & \\quad M\\lesssim p_T \\,, \\vspace{3mm}\\\\ \\alpha_s \\,g_\\chi^2\\, g_q^2\\, \\frac{p_T^2}{M^4} &\\quad  M\\gtrsim p_T \\,, \\\\ \\end{cases} \\label{eq:colliderqual} \\ee where $\\alpha_s$ is the QCD coupling and $p_T$ is the transverse momentum of the jet, which is typically $\\sim\\mathcal{O}(100)\\,\\gev$ at the Tevatron\\footnote{Note that (\\ref{eq:colliderqual}) is only qualitative in nature.  The limits are correct for mediator masses well above and below the $p_T$ of the jet.}.  Thus, for the heavy mediator case the (partonic) production cross section at the Tevatron, where $p_T\\sim 100\\ \\gev$, is $\\mathcal{O}(1000)$ times larger than the direct detection cross section for $\\mu \\sim 1$~GeV when the DM is heavier than the nucleon mass.  The CDF mono-jet search~\\cite{CDFmonojet} analysed $\\sim1\\ {\\rm fb}^{-1}$ and saw no significant discrepancy from the SM, thus limiting the DM + mono-jet production cross section to be smaller than $\\sim 500\\ {\\rm fb}$.  Due to the factor of 1000 mentioned above, this will translate to bounds in the neighborhood of $0.5$ fb in direct detection experiments, the exact bound at direct detection experiments will depend upon the details of the parton density functions relating the partonic cross section of (\\ref{eq:colliderqual}) to the actual CDF mono-jet bound.  This is to be compared with direct detection current searches. Null results from experiments such as CDMS~\\cite{Ahmed:2009zw}, XENON\\cite{Angle:2007uj,Aprile:2010um} and others, place strong constraints on the cross section of DM to recoil from a nucleus, $\\sigma \\lesssim 10^{-3}-10^{-4}~{\\rm fb}$ for a 10-100 GeV WIMP scattering elastically through a spin independent (SI) interaction.  Thus, for this situation it seems that direct detection has greater reach.  However, due to the threshold to detect a DM recoil in these experiments there is a DM mass below which these experiments are no longer sensitive, typically this lower bound is $m_\\chi\\sim 5-10\\ \\gev$, there is no such threshold in collider searches.  Furthermore, the DAMA collaboration~\\cite{Bernabei:2008yi} have observed a signal consistent with DM scattering from NaI % which is inconsistent with bounds on a standard WIMP from CDMS and other experiments. This has motivated the introduction of non-standard DM scenarios that can make these seemingly discrepant results consistent.  The cross sections necessary to explain DAMA are considerably larger than $10^{-3} {\\rm fb}$ and may allow these scenarios to be probed directly at the Tevatron, due to the increase in cross section described above. Another possibility that has been motivated both by DAMA and the recent CoGeNT~\\cite{Aalseth:2010vx} excess is that dark matter is light, below about 10 GeV, and is thus transfers small momenta to nuclei giving a signal near threshold. The Tevatron will place a strong bound for dark matter particles below 5 GeV. Finally, spin-dependent (SD) WIMP-nucleus scatterings are not coherent and therefore are not enhanced by an $A^2$ factor.  Typical bounds on a SD WIMP-proton scatter from direct detection are $\\sim 1$ fb , and will be severely impacted by the mono-jet bounds presented here.  We will begin our discussion with a model independent operator analysis, corresponding to very heavy mediation particles (such as a heavy $Z'$ or squarks). In Section~\\ref{sec:Ops} we will introduce some representative four fermion operators supressed by a cutoff scale. We will then place limits on the strengths of these operators from the Tevatron mono-jet search. In Section~\\ref{sec-DD} we will translate the Tevatron bounds to limits on direct detection cross section for different dark matter scenarios. In Section~\\ref{sec:lightmediator} we move on to introduce lighter mediators that are kinematically accesible at the Tevatron and find that these can either slightly enhance or severely weaken the Tevatron bounds. In Section~\\ref{sec-discussion} we will discuss possible enhancements to the Tevatron dark matter search using the mono-jet $p_T$ spectrum, and conclude. ",
        "conclusions": "\\label{sec-discussion}  \\begin{figure}[t] \\begin{center} \\includegraphics[width=0.48\\textwidth]{Puu10GeV.eps} \\vspace{2mm} \\includegraphics[width=0.48\\textwidth]{PCalcMad.eps}\\\\ (a)\\qquad\\qquad\\qquad\\qquad\\qquad\\qquad\\qquad\\qquad\\qquad\\qquad(b) \\caption{(a) Comparisons of the shapes of the signal, the SM background and CDF measured events. The SM predictions are shown in the green and the CDF observed data are shown in red. (b) Comparisons of simulated signal events from two different Monte-Carlo tools and for the parton and the particle levels. The cutoff $\\Lambda \\equiv M/\\sqrt{g_\\chi g_q}$ is chosen to be 1 TeV.} \\label{fig-spectra} \\end{center} \\end{figure}   It is worthwhile to consider possible improvements to the dark matter search at the Tevatron, and in the future at the LHC. Here we placed bounds on dark matter using only the total rate of mono-jet signal events above a certain $p_T$ cut. An analysis that takes the spectrum shape into account may yield more powerful bounds. We show the spectral shape of the signal compared to the background in Figure~\\ref{fig-spectra}(a). We find that the signal spectrum is somewhat harder than the background, especially when the messenger mass is much higher than the dark matter mass. We find that including showering, hadronization (using Pythia~\\cite{Sjostrand:2006za}) and a detector simulation (PGS~\\cite{PGS}) does not change the signal shape significantly, particularly above 100 GeV, as is shown in Figure~\\ref{fig-spectra}(b). This may allow us to place tighter constraints using a multi-bin analysis as compared with a simple counting experiment, since signal predicts more deviations in high $p_T$ bins. However, this would require knowledge of the theoretical uncertainty on a bin-by-bin basis which is not presently available.  Furthermore, a bound may be extracted from mono-photon events.   In this work we show that the Tevatron mono-jet search places competitive bounds on dark matter-nucleus cross sections relevant for direct detection experiments. In particular, the Tevatron limits are the current world-best for light dark matter, below a mass of 5 GeV. The Tevatron also sets the best limit spin dependent dark matter scattering. Various models built to explain the DAMA modulation signal such as inelastic and exothermic dark matter are also constrained by current Tevatron searches.  In addition to considering dark matter that couples to quarks via contact interactions we have taken the possibility of light mediators, as motivated by cosmic ray excesses~\\cite{ArkaniHamed:2008qn} into account. We find that the introduction of a light mediator of mass $\\lesssim 10$ GeV alleviates the mono-jet bounds completely for most cases.  This leads to an interesting conclusion - if a direct dark matter signal is established in a region that is in conflict with collider bounds, a new light state should be introduced to reconcile the data.  The Tevatron is unable to search for DM above a few 100 GeV due to kinematics, an upper bound that will be raised at the LHC. The current powerful bounds using only 1 fb$^{-1}$ motivates a dedicated analysis using more Tevatron data as well as future analyses at the LHC.     \\vspace{0.5cm} \\textbf{Note added:} While in the final stages of this work Ref.~\\cite{Goodman:2010yf} appeared on the arXiv.  They address similar issues, focussing on Majorana DM, but do not consider light mediators.  \\subsection*{Acknowledgements} We would like to thank John Campbell, Walter Giele, Joachim Kopp, Adam Martin and Neal Weiner for useful discussions.  We are grateful to Joachim Kopp for providing the contours used in Figure~\\ref{fig-iDM-limits}.  PJF would like to thank the CCCP at NYU, and the UC Davis HEFTI workshop on Light Dark Matter for hospitality while part of this work was undertaken. Fermilab is operated by Fermi Research Alliance, LLC, under Contract DE-AC02-07CH11359 with the United States Department of Energy.  \\appendix"
    },
    "1005/1005.1792_arXiv.txt": {
        "abstract": "In Ruderman \\& Sutherland (RS75) model, the normal neutron stars as pulsars bear a severe problem, namely the binding energy problem that both ions (e.g., ${}_{26}^{56}$Fe) and electrons on normal neutron star surface can be pulled out freely by the unipolar generator induced electric field so that sparking on polar cap can hardly occur. This problem could be solved within the Partially Screened Gap (PSG) model in the regime of neutron stars. However, in this paper we extensively study this problem in a bare strange quark star (BSS) model. We find that the huge potential barrier built by the electric field in the vacuum gap above polar cap could usually prevent electrons from streaming into the magnetosphere unless the electric potential of a pulsar is sufficiently lower than that at infinite interstellar medium. Other processes, such as the diffusion and thermionic emission of electrons have also been included here. Our conclusions are as follows: both positive and negative particles on a BSS's surface would be bound strongly enough to form a vacuum gap above its polar cap as long as the BSS is not charged (or not highly negative charged), and multi-accelerators could occur in a BSS's magnetosphere. Our results would be helpful to distinguish normal neutron stars and bare quark stars through pulsar's magnetospheric activities. ",
        "introduction": "Although pulsar-like stars have many different manifestations, they are populated mainly by rotation-powered radio pulsars. A lot of information about pulsar radiative process is inferred from the integrated and individual pulses, the sub-pulses, and even the micro-structures of radio pulses. Among the magnetospheric emission models, the user-friendly nature of Ruderman \\& Sutherland (1975; hereafter RS75) model is a virtue not shared by others~\\citep{Shukre92}.  In RS75 and its modified versions~\\citep[e.g.,][]{QL98}, a vacuum gap exists above polar cap of a pulsar, in which charged particles (electrons and positrons) are accelerated because of ${\\bf E \\cdot B} \\neq 0 $. These accelerated charged particles, moving along the curved magnetic field lines, radiate curvature or inverse-Compton-scattering-induced high energy photons which are converted to $e^\\pm$ while propagating in strong magnetic field. A follow-up breakdown of the vacuum gap produces secondary electron-positron pairs plasma that radiate coherent radio emission. These models with gap-sparking provide a good framework to analyze observational phenomena, especially the drifting~\\citep{DC68,DR99,VJ99} and bi-drifting~\\citep{QLZXW2004} sub-pulses.  However, the RS75-like vacuum gap models work only in strict conditions: strong magnetic field and low temperature on surface of pulsars~\\citep[e.g.,][]{GH08,ML07}. The necessary binding energy of positive ions (e.g., ${}_{26}^{56}$Fe) for RS75 model to work should be higher than $\\sim 10$ keV, while calculations showed that the cohesive energy of ${}_{26}^{56}$Fe at the neutron star surface is $<1$ keV~\\citep{FLRSHM77,L01}. This binding energy problem could be solved within a partially screened inner gap model~\\citep{GM03,GM06,MG09} for normal neutron stars. Alternatively, it is noted that the binding energy could be sufficiently high if pulsars are bare strange quark stars~\\citep{XQ98,XQZ99,XZQ01} although strange stars were previously supposed to exist with crusts~\\citep{AFO86}. Certainly, it is very meaningful in the elementary strong interaction between quarks and the phases of cold quark matter that the binding energy problem could be solved by bare quark stars as pulsars~\\citep{X09,X10}.  Though the ideas of solving the binding energy problem in BSS model were presented and discussed in some literatures, there is no comprehensive study with quantitative calculations up to now. In this paper, we are going to investigate the BSS model in quantitative details and show the physical picture of binding of particles on BSS's surface. Our research results are that multi-accelerators could occur above the polar cap for (and only for) the curvature-radiation-induced (CR-induced) sparking normal pulsars (NPs), but for other cases, such as resonant inverse-Compton-scattering-induced (ICS-induced) sparking NPs and both CR-induced and ICS-induced millisecond pulsars (MSPs), particles on surface of BSSs are bound strongly enough to form vacuum gap and RS75-like models work well if pulsars are BSSs. ",
        "conclusions": "In RS75 model, the binding energy problem is one of the most serious problems in the normal neutron star model of pulsars. Arons and Scharlemann (1979) developed an alternative model, the space-charge limited flow (SCLF) model, in which the particles, both iron ions and electrons can be pulled out freely, and form a steady flow~\\citep{AS79}. In this SCLF model, the drifting sub-pulse phenomenon which has been commonly observed in pulsars can rarely be reproduced. The prerequisite for understanding this phenomenon could be the existence of a vacuum gap.  In a very special case, through our calculations, we find that there is a new physical scenario for CR-induced sparking of normal pulsars (NPs) that free flow and vacuum gap may coexist above the polar cap. But in other cases, such as ICS-induced sparking of NPs and millisecond pulsars (MSPs), only vacuum gap exists. In general, if a pulsar is not highly negatively charged~\\citep{XCQ06}, vacuum gap survives at polar cap as well. One limitation is that our calculation is based on one-dimensional approximation and it might fail in some cases of MSPs. As far as we find, it is very difficult to deal with the high-dimensional cases. The one-dimensional approximation provides a good understanding of the geometry of polar cap of BSSs. In conclusion, the binding energy problem could be solved completely in the BSS model of pulsar as long as BSSs are neutral (or not highly negative charged), and the structure of polar cap of BSSs are very different with respect to that of NSs. Detailed information about the geometry of BSS's polar cap is given in Table~\\ref{AG}. \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{The accelerators above polar caps of BSSs. \\label{AG}} \\begin{tabular}{@{}llllll@{}} \\hline &\\multicolumn{2}{c}{[0,~$\\theta_{\\rm 0}$]$^{\\dag}$} &\\multicolumn{2}{c}{[$\\theta_{\\rm 0}$,~$\\theta_{\\rm A}$]} & \\\\ \\cline{2-5}\\\\ &CR &ICS &CR &ICS & \\\\ \\hline &SCLF  &VG &VG  &VG &$\\bf \\Omega \\cdot B > 0$ \\\\ NPs  &VG &VG$^{\\ddag}$ &SCLF &VG$^{\\ddag}$& $\\bf \\Omega \\cdot B < 0$ \\\\ &VG   &VG  &VG    &VG &$\\bf \\Omega \\cdot B >0 $ \\\\ MSPs &VG$^{\\ddag}$ &VG &VG$^{\\ddag}$ &VG &$\\bf \\Omega \\cdot B < 0$\\\\ \\hline \\end{tabular}\\\\ $^{\\dag}${$\\theta_{0}$ represents $\\theta_{\\rm 0,B}$ while choosing $\\phi_{\\rm B}=0$  and  $\\theta_{\\rm 0,C}$ while choosing  $\\phi_{\\rm C}=0$.}\\\\ $^{\\ddag}${for such cases, $\\theta_{0}$ $>$~$\\theta_{\\rm A}$, which represents the structure of the whole polar cap region.} \\end{minipage} \\end{table*} A more interesting region from pole to equator may locate between that polar angle where the total energy of electron equals the potential barrier and the polar angle of the foot of zero potential magnetic field line (i.e., [$\\theta_{0, \\rm C},\\theta_{\\rm C}$] or [$\\theta_{0, \\rm B},\\theta_{\\rm B}$], see Fig.~\\ref{PEC}) for CR-induced sparking NPs. After the birth of a NP, a vacuum gap exists at this region. When sparking starts, the potential in vacuum gap drops rapidly due to screen by electron-positron pairs and may become lower than that at the surface, namely $V_{\\rm i}(\\theta)$. As a result, the sparking converts vacuum gap to free flow at this region until the sparking ends, i.e., at [$\\theta_{0, \\rm C},\\theta_{\\rm C}$] or [$\\theta_{0, \\rm B},\\theta_{\\rm B}$], vacuum gap and free flow work alternately. This argument may have profound implications for us to distinguish neutron stars and quark stars by pulsar's magnetospheric activities (e.g., the diversity pulse profiles).  Another issue to be discussed is about the drifting rate of subpulses when we use the height of pure vacuum gap in this work. The natural explanation of the drifting subpulse phenomena in vacuum gap is due to $\\bf E \\times B$. Unfortunately, these theoretical calculations gave higher drifting rate with respect to observations~\\citep[e.g.,][]{RS75, DR99, DR01, GM03, GMZ06}. Since it has been observed~\\citep{DC68}, the drifting subpulse phenomenon remains unclear which has been widely regarded as one of the most critical and potentially insightful aspects of pulsar emission~\\citep{DR01}. The PSG mechanism~\\citep[e.g.,][]{GM03, GM06, GMZ06} could be a way to understand lower drifting rates observed, but some complexities still exist which make the underlying physics of drifting subpulses keep complicated and far from knowing clearly. (1) in principle the drifting velocity of subpulses is the ratio of the drifting distance to the duration, while the expected velocity predicted by $\\bf E \\times B$ is only for electrons in separated emission units, namely the plasma filaments. These two velocities would not be the same if the plasma filaments may stop after sparking. When sparking starts, the electric field in the vacuum gap vanishes due to screen by plasmas; while sparking ends, the electric field appears again. Thus, the calculated drifting velocity with $\\bf E \\times B$ could be higher than that of observations. (2) the so-called aliasing effect: as one observes subpulses only once every rotation period, we can hardly determine their actual speed. The main obstacles in the aliasing problem are the under sampling of subpulse motion and our inability to distinguish between subpulses especially when the differences between subpulses formed by various subbeams are smaller than the fluctuations in subpulses from one single subbeam~\\citep{LSRR03}. Anyway, detailed studies are very necessary in the future works.  We assume that the potential energy related to Eq.~\\ref{UGP}, $eV_{\\rm i}$, to be the constant, $\\phi_{0}$, in Eq.~\\ref{FD}. This assumption could be reasonable. For an uniformly magnetized, rotating conductor sphere, the unipolar generator will induce an electric field which is a function of polar angle, as described in Eq.~\\ref{UGP}. In the case of $\\bf \\Omega \\cdot B>0$ (Fig.~\\ref{antipulsar}), the potential energy of electron is highest at the polar region which means that those electrons there could be easier to escape. Alternatively, this conclusion could be quantitatively understood as following: because of Lorentz force inside a star, more electrons locate at the polar region so that the Fermi energy of electron is higher there and easier to escape into magnetosphere."
    },
    "1005/1005.1955_arXiv.txt": {
        "abstract": "The Megamaser Cosmology Project (MCP) aims to determine $H_0$ by measuring angular-diameter distances to galaxies in the Hubble flow using observations of water vapor megamasers in the circumnuclear accretion disks of active galaxies.  The technique is based only on geometry and determines $H_0$ in one step, independent of standard candles and the extragalactic distance ladder. In Paper I we presented a VLBI map of the maser emission from the Seyfert 2 galaxy UGC 3789. The map reveals an edge-on, sub-parsec disk in Keplerian rotation, analogous to the megamaser disk in NGC 4258. Here we present 3.2 years of monthly GBT observations of the megamaser disk in UGC 3789.  We use these observations to measure the centripetal accelerations of both the systemic and high-velocity maser components.  The measured accelerations suggest that maser emission lines near the systemic velocity originate on the front side of the accretion disk, primarily from segments of two narrow rings.  Adopting a two-ring model for the systemic features, we determine the angular-diameter distance to UGC 3789 to be 49.9 $\\pm$ 7.0 Mpc. This is the most accurate geometric distance yet obtained to a galaxy in the Hubble flow. Based on this distance, we determine $H_0$ = 69 $\\pm$ 11 km s$^{-1}$ Mpc$^{-1}$.  We also measure the mass of the central black hole to be 1.09 $\\times$ 10$^7$ $\\solmass$ $\\pm 14\\%$.  With additional observations the uncertainty in the distance to this galaxy can be reduced to under 10\\%.  Observations of megamaser disks in other galaxies will further reduce the uncertainty in $H_0$ as measured by the MCP. ",
        "introduction": "As a complement to observations of the cosmic microwave background (CMB), an independent measurement of the Hubble Constant, $H_0$, within 3\\% rms would place a valuable constraint on the dark energy equation-of-state parameter, $w$ (e.g. Hu 2005, Olling 2007). The most widely accepted value for $H_0$ is 72 $\\pm$ 7 km s$^{-1}$ Mpc$^{-1}$ (Freedman et al. 2001), based on the ``extragalactic distance ladder'' and using the Large Magellanic Cloud (LMC) to calibrate Cepheid variables which are treated as standard candles.  The Cepheid metallicity correction has been controversial, as highlighted by Sandage et al. (2006), who used similar methods but a different metallicity correction to obtain $H_0$ = 62 $\\pm$ 6 km s$^{-1}$ Mpc$^{-1}$.  If measured at the $\\lesssim$ 3\\% level, a value of $H_0$ = 72 km s$^{-1}$ Mpc$^{-1}$ would be consistent with the cosmological constant as an explanation of dark energy ($w$ = -1), whereas $H_0$ = 62 km s$^{-1}$ Mpc$^{-1}$ would challenge that model (e.g. Figure 14 of Riess et al., 2009).  Macri et al. (2006) used NGC 4258 rather than the LMC to calibrate Cepheids for the extragalactic distance scale and measured $H_0$ = 74 $\\pm$ 7 km s$^{-1}$ Mpc$^{-1}$.  Riess et al. (2009) subsequently refined the measurement, also using Cepheids in NGC 4258 to calibrate those in selected SN Ia host galaxies, and they determined $H_0$ = 74.2 $\\pm$ 3.6 km s$^{-1}$ Mpc$^{-1}$ (5\\%).  Riess et al. are able to reduce the systematic uncertainty in $H_0$ compared to the LMC-based results because the Cepheids in NGC 4258 are more similar in metallicity and period to those in the galaxies being calibrated.  NGC 4258 is a good foundation for the extragalactic distance scale because its distance (7.2 Mpc $\\pm$ 7\\%) is known to high accuracy from observations of circumnuclear water megamasers in its AGN accretion disk (Herrnstein et al. 1999). NGC 4258 is too close, though, for determining $H_0$ directly because its recession velocity may be highly influenced by its peculiar motion.  Lo (2005) gives a review of megamasers, including a discussion of the maser distance technique used in NGC 4258.  The Megamaser Cosmology Project (MCP) aims to determine $H_0$ by measuring angular-diameter distances to galaxies in the Hubble flow, $\\sim$ 50 -- 200 Mpc distant, using the maser technique.  Being independent of Cepheids and the CMB, the MCP will provide a valuable check of these other methods for measuring $H_0$. The project is a multi-year effort that begins with surveys for new megamaser disks.  Measuring a distance to each galaxy detected with a megamaser disk requires high-fidelity VLBI imaging, which we pursue with the High Sensitivity Array (VLBA+GBT+VLA+Effelsberg), and spectral-line monitoring to measure the centripetal accelerations of individual lines in the maser spectrum.  The MCP aims to measure maser distances with $\\sim$ 10\\% or better accuracy to each of about 10 galaxies, thereby determining H$_0$ to $\\sim$ 3\\%.  In a survey with the Green Bank Telescope, Braatz \\& Gugliucci (2008; hereafter BG08) discovered water maser emission from the Seyfert 2 galaxy UGC 3789.  The maser spectrum has three sets of Doppler components spaced nearly symmetrically at and around the galaxy's systemic recession velocity (Figure 1). This profile is characteristic of megamasers in an edge-on accretion disk.  The megamaser in UGC 3789 is similar to that in NGC~4258, where ``high-velocity'' blue-shifted and red-shifted maser components arise from the tangent points of the disk rotating toward and away from the observer, and components near the systemic recession velocity originate on the near side of the disk, along the line of sight to the central black hole.  Rotation velocities indicated by the high-velocity maser lines in UGC 3789 are $\\sim$ 400--800 km s$^{-1}$.  By monitoring the maser spectrum roughly monthly over 6 months, BG08 measured secular, redward drifts in radial velocity of 2 -- 8 km s$^{-1}$ yr$^{-1}$ among the systemic group of maser components, while the red- and blue-shifted features remained relatively fixed in velocity.  The radial velocity drift is interpreted as centripetal acceleration in the maser gas as it orbits the central supermassive black hole.  The spectral profile and the measured velocity drifts demonstrate that the maser disk in UGC 3789 is suitable for measuring an angular-diameter distance to the galaxy, using the method pioneered by Herrnstein et al. (1999).  In Paper I from the Megamaser Cosmology Project (Reid et al. 2009), we presented a VLBI image of the maser disk in UGC 3789.  The masers reveal nuclear gas in an edge-on disk ranging in radii from 0.08 to 0.30 pc.  The high-velocity lines trace a Keplerian rotation curve with high precision. While the high-velocity lines indicate relatively little disk warping, the systemic features show some structure that suggests the front side of the disk may be warped, tilted, or geometrically thick.  Here we present single-dish observations of the UGC 3789 maser made with the Green Bank Telescope over a period of 3.2 years.  Using these observations we refine the acceleration measurements made by BG08 and we present a model to estimate the distance to UGC 3789 using the maser observations. ",
        "conclusions": "The peculiar radial velocity of UGC 3789 relative to the CMB is -151 $\\pm$ 163 km s$^{-1}$ (Masters, private communication), using a flow model based on the SFI++ sample (Masters et al. 2006; Springob et al. 2007).  Here peculiar velocity is defined as $v_{pec} = v_{rec} - H_0 D$, where v$_{rec}$ is the galaxy's observed recession velocity.  Writing the UGC 3789 maser radial velocity in the CMB frame ($v_{CMB} \\simeq v_{LSR}$ + 60 km s$^{-1}$) and correcting for the peculiar velocity, we obtain a relativistic, recessional flow velocity of 3481 $\\pm$ 163 km s$^{-1}$.  Applying a standard $\\Lambda$CDM model with $\\Omega_m$ = 0.26, we therefore determine $H_0$ = 69 $\\pm$ 11 km s$^{-1}$ Mpc$^{-1}$.  The main source of uncertainty in the distance to UGC 3789 comes from the measurement of the orbital curvature parameter $\\Omega$.  The uncertainty in the acceleration is also a significant contributor, while the contribution from the Keplerian rotation constant is negligible.  Therefore we can best reduce the total uncertainty in the distance to UGC 3789 by obtaining a more sensitive VLBI map. We are currently processing additional VLBI observations of UGC 3789, which we can average to increase the map sensitivity.  Improving the measurement of the acceleration profile requires more sensitive monitoring observations.  Recognizing this need, we increased the duration of GBT observations after May 2007. These efforts should reduce the overall distance uncertainty to $<$ 10\\%, which would determine $H_0$ with accuracy comparable to the Hubble Key Project, but with entirely independent data and methods.  Ultimately, to reach an uncertainty in $H_0$ of $\\lesssim$ 3\\% from the Megamaser Cosmology Project, we require distances to at least 10 galaxies comparable to UGC 3789.  We are currently pursuing observations of the megamaser galaxies Mrk 1419, NGC 6264, and NGC 6323, which could eventually lead to comparable distance measurements.  There are another $\\sim$ 6 maser disk galaxies known that may also be suitable for distance measurements, but the quality of these (i.e. brightness of the masers, richness of the spectrum, and velocity coverage) is not as good.  Surveys for additional maser galaxies therefore remain an important part of the MCP."
    },
    "1005/1005.1487_arXiv.txt": {
        "abstract": "{The Photodetector Array Camera and Spectrometer (PACS) is one of the three science instruments on ESA's far infrared and submillimetre observatory. It employs two Ge:Ga photoconductor arrays (stressed and unstressed) with $16\\times25$ pixels, each, and two filled silicon bolometer arrays with $16\\times 32$ and $32\\times 64$ pixels, respectively, to perform integral-field spectroscopy and imaging photometry in the 60--210\\,$\\mu$m wavelength regime. In photometry mode, it simultaneously images two bands, 60--85$\\,\\mu$m or 85--125$\\,\\mu$m and 125--210$\\,\\mu$m, over a field of view of $\\sim 1.75'\\times 3.5'$, with close to Nyquist beam sampling in each band. In spectroscopy mode, it images a field of $47\\arcsec \\times 47\\arcsec$, resolved into $5\\times 5$ pixels, with an instantaneous spectral coverage of $\\sim 1500$~km/s and a spectral resolution of $\\sim 175$~km/s. We summarise the design of the instrument, describe observing modes, calibration, and data analysis methods, and present our current assessment of the in-orbit performance of the instrument based on the Performance Verification tests. PACS is fully operational, and the achieved performance is close to or better than the pre-launch predictions.} ",
        "introduction": "The PACS instrument was designed as a general-purpose science instrument covering the wavelength range $\\sim$60--210\\,$\\mu$m. It features both, a photometric multi-colour imaging mode, and an imaging spectrometer. Both instrument sections were designed with the goal of maximising the science return of the mission, given the constraints of the \\emph{Herschel} platform (telescope at $T \\approx 85\\ K$, diffraction limited for $\\lambda > 80\\,\\mu$m, limited real estate on the cryostat optical bench) and available FIR detector technology.  \\subsection {Photometer rationale}   Photometric colour diagnostics requires spectral bands with a relative bandwidth $\\Delta\\lambda/\\lambda < 0.5$. In coordination with the longer wavelength SPIRE bands, the PACS photometric bands have been defined as 60--85$\\,\\mu$m, 85--130$\\,\\mu$m, and 130--210$\\,\\mu$m, each spanning about half an octave in frequency.  A large fraction of the \\emph{Herschel} observing time will be spent on deep and/or large scale photometric surveys. For these, mapping efficiency is of the highest priority. Mapping efficiency is determined by both, the field of view of the instrument (in the diffraction-sampled case, the number of pixels) and the sensitivity per pixel. The PACS photometer was therefore designed around the largest detector arrays available without compromising sensitivity.  Simultaneous observation of several bands immediately multiplies observing efficiency.  By implementing two camera arrays, PACS can observe a field in two bands at a time.  Extracting very faint sources from the bright telescope background requires means to precisely flat-field the system responsivity on intermediate time-scales, as well as the use of spatial modulation techniques (chopping/nodding, scan-mapping) to move the signal frequency from ``DC'' into a domain above the 1/f ``knee'' of the system, including - most notably - the detectors.   \\subsection{Spectrometer rationale}  Key spectroscopic observations, particularly of extragalactic sources, ask for the detection of faint spectral lines with medium resolution ($R \\sim 1500$).  The power emitted or absorbed by a single spectral line in the far-infrared is normally several orders of magnitudes lower than the power in the dust continuum over a typical photometric band. Sensitivity is thus the most important parameter for optimisation; with background-limited detector performance the best sensitivity is obtained if the spectrometer satisfies the following conditions: The detection bandwidth should not be greater than the resolution bandwidth, which in turn should be matched to the line width of the source, and, the line flux from the source must be detected with the highest possible efficiency in terms of system transmission, spatial and spectral multiplexing.  Again, subtraction of the high telescope background has to be achieved by appropriate spatial and/or spectral modulation techniques. ",
        "conclusions": "With the PACS instrument, we have -- for the first time -- introduced large, filled focal plane arrays, as well as integral-field spectroscopy with diffraction-limited resolution, in the far infrared.  Without any prior demonstration, several major, technological developments have found their first application with PACS in space.  The success of this path -- documented by the results in this volume -- should encourage our community to defend this more ``pioneering'' approach against trends towards an ``industrial'' (i.e., minimal-risk) approach, which will not allow us to take advantage of the latest, experimental developments, on which our scientific progress often depends."
    },
    "1005/1005.3357_arXiv.txt": {
        "abstract": "We report the result of VLBI observation of the giant radio galaxy J1313+696 (4C +69.15) at 2.3/8.4 GHz, only the core component of the giant radio galaxy was detected in the VLBI observation at the dual frequencies. The result shows a steep spectrum core with $\\alpha=-0.82$ ($S \\propto \\nu^{\\alpha}$) between 2.3 GHz and 8.4 GHz. The steep spectrum core may be a sign of renewed activity. Considering also the upper limit flux density of 2.0 mJy at 0.6 GHz from Konar et al. 2004 the core has a GHz-peaked spectrum, implying that the core is compact and absorbed. Further high resolution VLBI observations are needed to identify if the steep spectrum core is consisting of a core and steep spectrum jet. ",
        "introduction": "It's known that the large radio sources are classified into two types i.e. FRI and FRII-type according to the morphology of the radio sources. Among the large radio sources, the radio galaxies whose lobes span a (projected) distance of above 1 Mpc are called giant radio galaxies (GRGs), and the majority of them are FRII sources (Schoenmakers et al. 2001). A small number of the large extended sources consisting of a pair of double lobes have been called double-double radio galaxies (DDRGs), and the observed structures of DDRGs suggest recurrent or interrupted central activity as the origin of these sources (Schoenmakers et al. 2000a; Saikia et al. 2006). The large radio sources have been observed with telescopes or arrays at relatively low resolution but not well observed with the very long baseline interferometry (VLBI). The VLBI observation is able to identify the core of the large radio sources at high resolution. We have included the GRG J1313+696 in our sample of VLBI observation with the European VLBI Network (EVN) at dual frequency 2.3/8.4 GHz simultaneously. ",
        "conclusions": "J1313+696 (B1312+698, DA340, 4C +69.15, z=0.106), is an FRII type giant radio galaxy. In VLA observation at 1.4 GHz the emission from the core to the lobes forms a continuous bridge in SE-NW direction (Lara et al. 2001), the GMRT observation at 605 MHz shows a similar structure (Konar et al. 2004). But at 4.9 GHz only the core and the lobe extremes are detected with the VLA (Lara et al. 2001), and the core position is consistent with the position of the associated galaxy. The source was not classified as a DDRG according to the definition of Schoenmakers et al. (2000a) which an inner pair of edge-brightened lobes should be detected.  The VLBI images (Fig.~\\ref{fig1}, Fig.~\\ref{fig2}) of J1313+696 at 2.3/8.4 GHz show a point source. We checked the position of the point source in the VLBI maps with the VLA images at 1.6 and 5 GHz (Lara et al. 2001), confirmed that the VLBI point source is the core of J1313+696 in the VLA images. We collected the core flux densities of J1313+696 from our data and literature as shown in Fig.~\\ref{fig3}, the error bars are 1$\\sigma$ (per beam) from our data and the literature. The source flux densities in the VLBI observation at 2.3 GHz and 8.4 GHz are 7.5$\\pm$1.1 mJy and 2.6$\\pm$0.2 mJy respectively. It is 10.2 mJy at 1.4 GHz and 3.8 mJy at 4.9 GHz in Lara et al. (2001), and 6 mJy at 2.7 GHz in Saunders et al. (1987). It is 7 mJy at 1.4 GHz, 4.3 mJy at 4.9 GHz, and an upper limit of 2 mJy at 605 MHz in Konar et al. (2004). The spectral index between 2.3 and 8.4 GHz in the VLBI observation is $-0.82$ (we use $S \\propto \\nu^{\\alpha}$). The spectral index from the collected data above 1.4 GHz can be fitted linearly (Fig.~\\ref{fig3}) with $\\alpha=-0.78\\pm0.03$. It is interesting that this value is consistent with the spectral indices of the east, west lobe of J1313+696 (Schoenmakers et al. 2000b). However, because an upper limit 2 mJy at 605 MHz was estimated by Konar et al. (2004) the integrated spectrum in Fig.~\\ref{fig3} can also be an inverted spectrum which peaked around 1.4 GHz. We note that the 1.4 GHz core flux (10.2 mJy) from Lara et al. (2001), might be overestimated due to contamination of diffuse emission. As Lara et al. mentioned, for this core the ratio peak/total $<0.8$. An ideal point source should have the ratio peak/total $\\sim$1. The core is sitting on top of a diffuse emission, as is evident from the map in Lara et al.. Whereas, the measurement of Konar et al. (2004) has less contamination as they re-mapped the field with lower uv-cutoff to lose the diffuse flux and get the core flux as correct as possible.  Konar et al. (2004) have discussed the steep-spectrum cores (SSCs, $\\alpha_{core} < -0.5$) of GRGs, in their sample at least 3 out of 17 sources show SSCs. In Lara et al. (2004) sample a compact radio core was detected at 4.9 GHz with VLA in all sources, with an average spectral index of $0.07\\pm0.41$ for FRII and $-0.24\\pm0.52$ for FRI sources; hence FRIs seem to have more SSCs than FRIIs according to the averaged spectral indices. However, Lara et al. (2004) explained that the core spectral index in FRIs suffers from more contamination from the steeper jet emission than in FRIIs. In Konar et al. (2004) sample, 3 FRII-type sources have been suggested to have SSCs besides J1313+696. They suggested that SSC is preferentially occurred in giant radio galaxies, and SSC is related to the interrupted activity of GRGs. Schoenmakers et al. (2000b) found a polarized component of J1313+696 near but not at the core position, they suggested it is a jet component.  The VLBI images of J1313+696 show a point source, not as expected to show a core-jet, but reveal a steep spectrum core which may imply a renewed jet in the core. Considering the upper limit flux density at 605 MHz the core of the J1313+696 has a GHz-peaked spectrum (GPS), indicating the core is very compact and absorbed at lower frequency. We estimated an upper limit of the core size at 8.4 GHz with the beam size of our VLBI map, it is 3.8 pc (in $H_0=71~km s^{-1} Mpc^{-1}$, $\\Omega_m$=0.27, and $\\Omega_{vac}$=0.73 cosmology). Another example is the DDRG B1834+620, its core shows a steep spectrum at higher frequencies with spectral index the same as that of lobes, and the core spectrum is inverted at lower frequencies, showing a GPS shape (Schoenmakers et al. 2000c). Konar et al. (2008) found a GPS core in the GRG J1155+4029 with GMRT and VLA observations."
    },
    "1005/1005.3816_arXiv.txt": {
        "abstract": "{ We present an early broad-brush analysis of \\emph{Herschel}/PACS observations of star-forming galaxies in 8 galaxy clusters drawn from our survey of 30 clusters at $z{\\simeq}0.2$.  We define a complete sample of $192$ spectroscopically confirmed cluster members down to $L_{\\rm TIR}{>}3{\\times}10^{10}{\\Lsol}$ and $L_K{>}0.25{\\Lsol}$.  The average $K$-band and bolometric infrared luminosities of these galaxies both fade by a factor of ${\\sim}2$ from clustercentric radii of ${\\sim}2r_{200}$ to ${\\sim}0.5r_{200}$, indicating that as galaxies enter the clusters ongoing star-formation stops first in the most massive galaxies, and that the specific star-formation rate (SSFR) is conserved.  On smaller scales the average SSFR jumps by ${\\sim}25\\%$, suggesting that in cluster cores processes including ram pressure stripping may trigger a final episode of star-formation that presumably exhausts the remaining gas.  This picture is consistent with our comparison of the \\emph{Herschel}-detected cluster members with the cluster mass distributions, as measured in our previous weak-lensing study of these clusters.  For example, the spatial distribution of the \\emph{Herschel} sources is positively correlated with the structures in the weak-lensing mass maps at ${\\sim}5{\\sigma}$ significance, with the strongest signal seen at intermediate group-like densities.  The strong dependence of the total cluster IR luminosity on cluster mass -- $L_{\\rm TIR}{\\propto}M_{\\rm virial}^2$ -- is also consistent with accretion of galaxies and groups of galaxies (i.e.\\ the substructure mass function) driving the cluster IR luminosity.  The most surprising result is that roughly half of the \\emph{Herschel}-detected cluster members have redder $S_{100}/S_{24}$ flux ratios than expected, based on the Rieke et al.\\ models.  On average cluster members are redder than non-members, and the fraction of red galaxies increases towards the cluster centers, both of which indicate that these colors are not attributable to systematic photometric errors.  Our future goals include to intepret physically these red galaxies, and to exploit this unique large sample of clusters with unprecedented multi-wavelength observations to measure the cluster-cluster scatter in S0 progenitor populations, and to intepret that scatter in the context of the hierarchical assembly of clusters. } ",
        "introduction": "Lenticular galaxies (hereafter S0s) are mainly found in the cores of galaxy clusters at low redshift (e.g.\\ Dressler et al.\\ 1997; Smith et al.\\ 2005; Postman et al.\\ 2005).  There is a broad consensus that they are the descendants of gas rich spiral galaxies that have been accreted from the surrounding filamentary structure.  However the physics of how spirals are transformed into S0s remains largely unconstrained, with numerous ``S0 progenitor'' populations (e.g.\\ Moran et al.\\ 2006; Poggianti et al.\\ 2000; Geach et al.\\ 2006; Haines et al.\\ 2009a -- hereafter H09a) and physical processes (e.g.\\ Gunn \\& Gott 1972; Moore et al.\\ 1999) discussed in the literature.  The broad range of cluster-centric radii at which various S0 progenitors are found reflects the fact that different physical processes act in different environments, for example ram pressure stripping is more effective closer to cluster centers where the intracluster medium (ICM) is denser, and galaxy-galaxy merging is more effective in galaxy groups that are falling into the cluster than in the cluster cores.  Moreover, the observational signatures of S0 progenitors are diverse, ranging in wavelength from ultraviolet (UV) emission from A stars in galaxies whose star-formation (SF) has been recently quenched, through optical spectral features including Balmer absorption lines, to mid/far-infrared (IR) emission from dust heated by SF (e.g.\\ Moran et al.\\ 2007; Poggianti et al.\\ 2000; Haines et al.\\ 2009b).  Mid- and far-IR properties of cluster galaxies have been studied previously with \\emph{IRAS} (e.g.\\ Leggett et al.\\ 1987; Doyon \\& Joseph 1989), \\emph{ISO} (see Metcalfe et al.\\ 2005 for a review), and \\emph{Spitzer} (e.g.\\ Geach et al.\\ 2006; Fadda et al.\\ 2008; Haines et al.\\ 2009a,b; Bai et al.\\ 2009).  A key result from these IR studies is that a significant fraction of the total SF in galaxy clusters is obscured by dust.  The inferred levels of SF naturally fit the hypothesis that bulge dominated S0s are descended from late-type spirals.  It has also been suggested that dusty S0 progenitors are more common in dynamically active, i.e.\\ merging, galaxy clusters than in so-called ``relaxed'' clusters (e.g.\\ Metcalfe et al.\\ 2005; Geach et al.\\ 2006; Miller et al.\\ 2006).  However it has thus far been difficult to test this idea robustly because the intrinsic scatter in levels of SF in clusters appears to be large, (as noted by Kodama et al.\\ 2004), and the sample sizes observed to date are small (i.e.\\ $\\ls2$) within any given redshift bin -- although see H09a for a recent counter-example.  We are therefore conducting a systematic wide-field survey of a large statistically well-defined sample of galaxy clusters in a narrow redshift slice at $z{\\simeq}0.2$, as part of the Local Cluster Substructure Survey (LoCuSS\\footnote{http://www.sr.bham.ac.uk/locuss/}).  Our goals are to compile a complete inventory of S0 progenitors using data from the far-UV to far-IR, and to relate these populations to the underlying gas physics and hierarchical structure of the host galaxy clusters. We aim to delineate the different physical processes responsible for galaxy transformation in clusters and their surrounding large scale structure, and thus constrain the amplitude of the different physical pathways from spiral to S0 morphology, and how these relate to the dynamical state of the clusters.  Our Open Time Key Programme observations with \\emph{Herschel} (Pilbratt et al., 2010), supplemented by existing \\emph{Spitzer} mid-IR observations provide the all-important measurements of the bolometric IR luminosity and mid/far-IR colors of dust-reddened/obscured S0 progenitors.  We assume ${\\Ho}{=}70{\\kms}$, ${\\Om}{=}0.3$, ${\\Ol}{=}0.7$.  In this cosmology $1{\\kpc}$ at $z{=}0.2$, subtends $0.3''$.  All cluster masses and radii relative to an over-density are derived from the weak-lensing analysis of Okabe et al.\\ (2010; hereafter Ok10). ",
        "conclusions": "We have presented an initial broad-brush analysis of \\emph{Herschel}/PACS observations of $25\\%$ of our sample of 30 galaxy clusters at $z{\\simeq}0.2$, and combined these data with our existing \\emph{Spitzer}, Subaru, \\emph{Chandra}, UKIRT, and MMT data.  The main analysis concentrates on a sample of 192 spectroscopically confirmed cluster members with $L_{\\rm TIR}{>}3{\\times}10^{10}L_\\odot$, $L_K{>}0.25L_K^\\star$, $R{<}1.5r_{200}$.  The average $K$-band luminosity of these galaxies fades by a factor of almost 2 from the cluster outskirts (${\\sim}1{-}2r_{200}$) to the cluster cores (${\\ls}0.5r_{200}$), although the average specific star-formation rate, as probed by ${\\langle}L_{\\rm TIR}{\\rangle}/{\\langle}L_K{\\rangle}$, is constant across most of this radial range (${\\sim}0.5{-}2r_{200}$), before jumping by 25\\% on smaller scales.  This suggests that as gas rich galaxies fall into the clusters (typically in groups -- Fig.~1) ongoing star-formation stops first in the most massive galaxies.  As galaxies reach the cluster cores physical processes that operate in high density environments, for example ram pressure stripping and harrassment, then appear to trigger a final episode of star-formation that presumably exhausts the remaining gas supply.  This picture is consistent with our comparison of the \\emph{Herschel}-detected cluster members with the cluster mass distributions, as probed by Okabe et al.'s (2010) weak-lensing analysis.  First, the spatial distribution of the \\emph{Herschel} sources is positively correlated with the structures in the weak-lensing mass maps at ${\\sim}5{\\sigma}$ significance, with the strongest signal seen at intermediate, group-like densities.  Second, the strong dependence of the total cluster IR luminosity on cluster mass ($L_{\\rm TIR}{\\propto}M_{\\rm virial}^2$) is consistent with accretion of galaxies and groups of galaxies driving the cluster IR luminosity.  This is because, assuming that IR galaxy mass-to-light ratios are independent of the cluster mass, the scaling relation can be understood as stemming from the $M^2$ dependence of the substructure mass function seen in theoretical models (e.g.\\ Taylor \\& Babul 2005).  Third, the IR luminosity density profiles of the clusters generally increase to large radii, with some clusters showing a peak at ${\\sim}r_{200}$.  This is qualitatively consistent with the enhanced star-formation rates seen in in-falling galaxy populations by Moran et al.\\ (2005).  The most surprising result is that roughly half of the \\emph{Herschel}-detected cluster galaxies have excess flux at $100{\\um}$ over that predicted from current SED models.  Cluster members are redder than non-members, and we find a shallow trend of increasing fraction of red IR galaxies towards the cluster centers, both of which suggest that this is a physical effect and not caused by systematic photometric uncertainties.  This result will be the focus of more detailed future investigation.  Finally, we note that, contrary to previous speculation in the literature, we do not find a strong relationship between cluster IR luminosity and cluster dynamical state.  Observations of the full sample will allow us to investigate this issue in more detail."
    },
    "1005/1005.3027_arXiv.txt": {
        "abstract": "We present near-infrared Ks-band photometry bracketing the secondary eclipse of the hot Jupiter TrES-2b using the Wide-field Infrared Camera on the Canada-France-Hawaii Telescope. We detect its thermal emission with an eclipse depth of \\FpOverFStarPercentAbstractTrESTwo$^{+\\FpOverFStarPercentAbstractPlusTrESTwo}_{-\\FpOverFStarPercentAbstractMinusTrESTwo}$\\% (\\XSigmaTrESTwo $\\sigma$). Our best-fit secondary eclipse is consistent with a circular orbit (a 3$\\sigma$ upper limit on the eccentricity, $e$, and argument or periastron, $\\omega$, of $|$$e$$\\cos$$\\omega$$|$ $<$ \\ECosOmegaAbsoluteThreeSigmaLimitTrESTwo), in agreement with mid-infrared detections of the secondary eclipse of this planet. A secondary eclipse of this depth corresponds to a day-side Ks-band brightness temperature of $T_B$ = \\TBrightTrESTwo$^{+\\TBrightPlusTrESTwo}_{-\\TBrightMinusTrESTwo}$ $K$. Our thermal emission measurement when combined with the thermal emission measurements using Spitzer/IRAC from O'Donovan and collaborators suggest that this planet exhibits relatively efficient day to night-side redistribution of heat and a near isothermal dayside atmospheric temperature structure, with a spectrum that is well approximated by a blackbody. It is unclear if the atmosphere of TrES-2b requires a temperature inversion; if it does it is likely due to chemical species other than TiO/VO as the atmosphere of TrES-2b is too cool to allow TiO/VO to remain in gaseous form. Our secondary eclipse has the smallest depth of any detected from the ground at around 2 $\\mu m$ to date. ",
        "introduction": "The first detection of the transit of an exoplanet in front of its parent star (\\citealt{Charb00}; \\citealt{Henry00}) opened a new avenue to determine the characteristics of these exotic worlds. For all but the most eccentric cases, approximately half-an-orbit after their transits these planets pass behind their star along our line of sight allowing their thermal flux to be measured in the infrared. The first detections of an exoplanet's thermal emission (\\citealt{Charb05}; \\citealt{Deming05}) came from observations in space with Spitzer using the Infrared Array Camera (IRAC; \\citealt{Fazio04}). Since then the vast majority of such measurements have been made using Spitzer at wavelengths longer than 3 $\\mu m$, and thus longwards of the blackbody peak of these ``hot'' exoplanets. Recent observations have extended secondary eclipse detections into the near-infrared; the first detection was from space with NICMOS on the Hubble Space Telescope (\\citealt{Swain09} at $\\sim$2 $\\mu m$). More recently, near-infrared detections have been achieved from the ground; the first of these detections include a $\\sim$6$\\sigma$ detection in K-band of TrES-3b using the William Herschel Telescope \\citep{deMooij09}, a $\\sim$4$\\sigma$ detection in z'-band of OGLE-TR-56b using Magellan and the Very Large Telescope (VLT; \\citealt{Sing09}), and a $\\sim$5$\\sigma$ detection at $\\sim$2.1 $\\mu m$ of CoRoT-1b also with the VLT \\citep{Gillon09}.  Thermal emission measurements in the near-infrared are crucial to our understanding of these planets' atmospheres, as they allow us to constrain hot Jupiters' thermal emission near their blackbody peaks. The combination of Spitzer/IRAC and near-infrared thermal emission measurements allows us to constrain the temperature-pressure profiles of these planets' atmospheres over a range of pressures \\citep{Fortney08}, better estimate the bolometric luminosity of these planets' dayside emission, and thus contributes to a more complete understanding of how these planets transport heat from the day to nightside at a variety of depths and pressures in their atmospheres \\citep{Barman08}.  The transiting hot Jupiter TrES-2b orbits a G0 V star with a period of $\\sim$2.47 $d$ \\citep{ODonovan06}. According to the \\citet{Fortney08} theory this places TrES-2b marginally in the hottest, mostly highly irradiated class (the pM-class) of hot Jupiters and close to the dividing line between this hottest class and the merely warm class of hot Jupiters (the pL-class). Thus TrES-2b could be a key object to refine the dividing line between these two classes, and indicate the physical cause of this demarcation, or reveal whether this divide even exists. Recently \\citet{ODonovan09} used Spitzer/IRAC to measure the depth of the secondary eclipse of TrES-2b in the four IRAC bands. Their best-fit eclipses are consistent with a circular orbit, and collectively they are able to place a 3$\\sigma$ limit on the eccentricity, $e$, and argument of periastron, $\\omega$, of $|$$e$cos$\\omega$$|$ $<$ 0.0036. Their best-fit eclipses at 3.6, 5.8 and 8.0 $\\mu m$ are well-fit by a blackbody. At 4.5 $\\mu m$ they detect excess emission, in agreement with the theory of several researchers (\\citealt{Fortney08, Burrows08}) that predicts such excess due to water emission, rather than absorption, at this wavelength due to a temperature inversion in the atmosphere. One-dimensional radiative-equilibrium models for hot Jupiter planets generally show that the atmospheric opacity is dominated by water vapor, which is especially high in the mid-infrared, but has prominent windows (the JHK bands) in the near infrared \\citep{Fortney08,Burrows08}. One can probe more deeply, to gas at higher pressure, in these opacity windows.  Models without temperature inversions feature strong emission in the JHK bands, since one sees down to the hotter gas.  Models with temperature inversions, since they feature a relatively hotter upper atmosphere and relatively cooler lower atmosphere, yield weaker emission in the near-IR (JHK), but stronger emission in the mid-infrared \\citep{Hubeny03,Fortney06}. Near-infrared thermal emission measurements should thus be useful to determine whether TrES-2b does or does not harbour a temperature inversion.  Owing to its high irradiation, with an incident flux of $\\sim$$1.1$$\\times$$10^{9}$ $erg$$s^{-1}$$cm^{-2}$, and favourable planet-to-star radius ratio ($R_{P}/R_{*}$$\\sim$0.13), we included TrES-2b in our program observing the secondary eclipses of some of the hottest of the hot Jupiters from the ground. Here we present Ks-band observations bracketing TrES-2b's secondary eclipse using the Wide-field InfraRed Camera (WIRCam) on the Canada-France-Hawaii Telescope (CFHT). We report a \\XSigmaTrESTwo $\\sigma$ detection of its thermal emission. ",
        "conclusions": "The depth of our best-fit secondary eclipse is \\FpOverFStarPercentAbstractTrESTwo$^{+\\FpOverFStarPercentAbstractPlusTrESTwo}_{-\\FpOverFStarPercentAbstractMinusTrESTwo}$\\%. The reduced $\\chi^{2}$ is \\ChiTrESTwo. Our best-fit secondary eclipse is consistent with a circular orbit; the offset from the expected eclipse center is: $t_{offset}$ = \\TOffsetTrESTwo$^{+\\TOffsetPlusTrESTwo}_{-\\TOffsetMinusTrESTwo}$ minutes (or at a phase of $\\phi$=\\PhaseAbstractTrESTwo$^{+\\PhaseAbstractPlusTrESTwo}_{-\\PhaseAbstractMinusTrESTwo}$). This corresponds to a limit on the eccentricity and argument of periastron of $e \\cos \\omega$ = \\ECosOmegaTrESTwo$^{+\\ECosOmegaPlusTrESTwo}_{-\\ECosOmegaMinusTrESTwo}$, or a 3$\\sigma$ limit of $|$$e$$\\cos$$\\omega$$|$ $<$ \\ECosOmegaAbsoluteThreeSigmaLimitTrESTwo). Our result is fully consistent with the more sensitive $e$cos$\\omega$ limits reported from the secondary eclipse detections at the four Spitzer/IRAC wavelengths \\citep{ODonovan09}. Thus our result bolsters the conclusion of \\citet{ODonovan09} that tidal damping of the orbital eccentricity is unlikely to be responsible for ``puffing up'' the radius of this exoplanet.  A secondary eclipse depth of \\FpOverFStarPercentAbstractTrESTwo$^{+\\FpOverFStarPercentAbstractPlusTrESTwo}_{-\\FpOverFStarPercentAbstractMinusTrESTwo}$\\% corresponds to a brightness temperature of $T_{B}$ = \\TBrightTrESTwo$^{+\\TBrightPlusTrESTwo}_{-\\TBrightMinusTrESTwo}$ $K$ in the Ks-band assuming the planet radiates as a blackbody, and adopting a stellar effective temperature of $T_{eff}$ = 5850 $\\pm$ 50 \\citep{Sozzetti07}. This compares to the equilibrium temperature of TrES-2b of $T_{eq}$$\\sim$1472 $K$ assuming isotropic reradiation, and a zero Bond albedo. Hot Jupiter thermal emission measurements allow joint constraints on the Bond albedo, $A_B$, and the efficiency of day to nightside redistribution of heat on these presumably tidally locked planets. The Bond albedo, $A_B$ is the fraction of the bolometric, incident stellar irradiation that is reflected by the planet's atmosphere. We parameterize the redistribution of dayside stellar radiation absorbed by the planet's atmosphere to the nightside by the reradiation factor, $f$, following the \\citet{Lopez07} definition. If we assume a Bond albedo near zero, consistent with observations of other hot Jupiters \\citep{Charbonneau99,Rowe08} and with model predictions \\citep{Burrows08b}, we find a reradiation factor of $f_{Ks}$ = \\fReradiationTrESTwo$^{+\\fReradiationPlusTrESTwo}_{-\\fReradiationMinusTrESTwo}$ from our Ks-band eclipse photometry only, indicative of relatively efficient advection of heat from the day-to-nightside at this wavelength. In comparison, the reradiation factor for an atmosphere that reradiates isotropically is $f$=$\\frac{1}{4}$, while $f$=$\\frac{1}{2}$ denotes redistribution and reradiation over the dayside face only.  \\begin{figure} \\centering \\includegraphics[scale=0.65,angle=270]{f8.eps} \\caption{\tThe 68.3\\% (1$\\sigma$; solid-line), 95.5\\% (2$\\sigma$; dashed-line) and 99.7\\% (3$\\sigma$; short dashed-line) $\\chi^2$ confidence regions on the reradiation factor, $f_{tot}$, and Bond albedo from the combination of our Ks-band point and the Spitzer/IRAC measurements \\citep{ODonovan09}.} \\label{FigBondReradiation} \\end{figure}  Our secondary eclipse depth, when combined with the secondary eclipse depths at the Spitzer/IRAC wavelengths from \\citet{ODonovan09}, is consistent with a range of Bond albedos, $A_B$, and efficiencies of the day to nightside redistribution on this presumably tidally locked planet (Figure \\ref{FigBondReradiation}). The best-fit total reradiation factor, $f_{tot}$, that results from a $\\chi^{2}$ analysis of all the eclipse depths for TrES-2b assuming a zero Bond albedo is $f_{tot}$ = \\fReradiationTrESTwoALL$^{+\\fReradiationPlusTrESTwoALL}_{-\\fReradiationMinusTrESTwoALL}$. Thus our Ks-band brightness temperature ($T_{B}$ = \\TBrightTrESTwo$^{+\\TBrightPlusTrESTwo}_{-\\TBrightMinusTrESTwo}$ $K$) and reradiation factor $f_{Ks}$=\\fReradiationTrESTwo$^{+\\fReradiationPlusTrESTwo}_{-\\fReradiationMinusTrESTwo}$, reveal an atmospheric layer that is similar to, and perhaps slightly hotter, than the atmospheric layers probed by longer wavelength Spitzer observations ($T_B$$\\sim$1500 $K$ from Spitzer/IRAC observations of TrES-2b [\\citealt{ODonovan09}]). The Ks-band is expected to be at a minimum in the water opacity \\citep{Fortney08,Burrows08}, and thus our Ks-band observations are expected to be able to see deep into the atmosphere of TrES-2b. Our observations suggest that the deep, high pressure atmosphere of TrES-2b displays a similar temperature -- perhaps a slightly warmer temperature -- to lower pressure regions.  Another way of parameterizing the level of day-to-nightside heat redistribution is calculating the percentage of the bolometric luminosity emitted by the planet's dayside, $L_{day}$, compared to the nightside emission, $L_{night}$. Measurements of the thermal emission of a hot Jupiter at its blackbody peak provide a valuable constraint on the bolometric luminosity of the planet's dayside emission, and by inference its nightside emission \\citep{Barman08}. From simple thermal equilibrium arguments if TrES-2b has a zero Bond albedo and it is in thermal equilibrium with its surroundings it should have a total bolometric luminosity of $L_{tot}$ = 7.7$\\times$10$^{-5}$$L_{\\odot}$. By integrating the luminosity per unit frequency of our best-fit blackbody model across a wide wavelength range we are able to calculate the percentage of the total luminosity reradiated by the dayside as $\\sim$\\DaysidePercentage\\% ($L_{day}$ = \\Lday$\\times$10$^{-5}$$L_{\\odot}$). The remainder, presumably, is advected via winds to the nightside.  \\begin{figure} \\centering \\includegraphics[scale=0.465,angle=270]{f9a.eps} \\includegraphics[scale=0.465,angle=270]{f9b.eps} \\caption{ Dayside planet-to-star flux ratios (top) and dayside flux at the planet's surface (bottom). The Ks-band point (black point; $\\sim$2.15 $\\mu m$) is our own, while the Spitzer/IRAC red points are from \\citet{ODonovan09}. Blackbody curves for isotropic reradiation ($f$=$\\frac{1}{4}$; $T_{eq}$$\\sim$1496 $K$; blue dashed-line) and for our best-fit reradiation factor ($f$=0.346; $T_{eq}$$\\sim$1622 $K$; grey dotted-line) are also plotted. We also plot one-dimensional, radiative transfer spectral models \\citep{Fortney06,Fortney08} for various reradiation factors and with and without TiO/VO. The models with TiO/VO include $f$=$\\frac{1}{4}$ (purple dotted line), $f$=0.31 (green dashed line), and $f$ =$\\frac{1}{2}$ (orange dotted line); only the last of the models has a temperature inversion. The model without TiO/VO features emission from the dayside only ($f$=$\\frac{1}{2}$; cyan dot-dashed line). The models on the top panel are divided by a stellar atmosphere model \\citep{Hauschildt99} of TrES-2 using the parameters from \\citet{Torres08} ($M_{*}$=0.98 $M_{\\odot}$, $R_{*}$=1.00 $R_{\\odot}$, $T_{eff}$=5850 $K$, and log $g$= 4.43). We plot the Ks-band WIRCam transmission curve (dotted black lines) and Spitzer/IRAC curves (solid red lines) inverted at arbitrary scale at the top of both panels. } \\label{FigModel} \\end{figure}   We compare the depth of our Ks-band eclipse and the Spitzer/IRAC eclipses \\citep{ODonovan09} to a series of planetary atmosphere models in Figure \\ref{FigModel}. This comparison is made quantitatively as well as qualitatively by integrating the models over the WIRCam Ks band-pass as well as the Spitzer/IRAC channels, and calculating the $\\chi^{2}$ of the thermal emission data compared to the models. We first plot blackbody models with an isotropic reradiation factor ($f$=$\\frac{1}{4}$; blue dotted-line) and that of our best-fit value ($f$=0.346; grey dotted-line) these models have dayside temperatures of $T_{day}$$\\sim$1496$K$ and $T_{day}$$\\sim$1622$K$, respectively. Both blackbody models provide reasonable fits to the data, although the latter model ($f$=0.346; $\\chi^{2}$=\\BlackbodyTwoChi) provides a definitively better fit than the former isothermal model ($f$=$\\frac{1}{4}$; $\\chi^2$=\\BlackbodyOneChi) as it better predicts our Ks-band emission and the Spitzer/IRAC 8.0 $\\mu m$ emission. This suggests that overall TrES-2b has a near-isothermal dayside temperature-pressure profile and is well-fit by a blackbody.  We thus also compare the data to a number of one-dimensional, radiative transfer, spectral models \\citep{Fortney06,Fortney08} with different reradiation factors that specifically include or exclude gaseous TiO and VO into the chemical equilibrium and opacity calculations. In these models when TiO and VO are present they act as absorbers at high altitude and lead to a hot stratosphere and a temperature inversion \\citep{Hubeny03}. However, if the temperature becomes too cool (TiO and VO start to condense at 1670 $K$ at 1 mbar [\\citealt{Fortney08}]), TiO and VO condense out and the models with and without TiO/VO are very similar. In the case of TrES-2b, for all the models we calculate, except our model that features dayside emission only ($f$=$\\frac{1}{2}$), they do not harbour temperature inversions because the atmospheres are slightly too cool and TiO/VO has condensed out of their stratospheres. We plot models with TiO/VO and reradiation factors of $f$=$\\frac{1}{4}$ (purple dotted line), $f$=0.31 (green dashed line), and $f$=$\\frac{1}{2}$ (orange dotted line), and without TiO/VO with a reradiation factor of $f$=$\\frac{1}{2}$ (cyan dot-dashed line). \\citet{ODonovan09} argued that TrES-2b experienced a temperature inversion due to the high 4.5 $\\mu m$ emission compared to the low 3.6 $\\mu m$ emission, which was predicted to be a sign of water and CO in emission, rather than absorption, in TrES-2b's presumably inverted atmosphere. We also find that our models without a temperature inversion have difficultly matching the Spitzer/IRAC 5.6 and 8.0 $\\mu m$ thermal emission ($\\chi^{2}$=\\FortneyTwoChi \\ for $f$=$\\frac{1}{4}$ with TiO/VO, $\\chi^{2}$=\\FortneyOneChi \\ for $f$=0.31 with TiO/VO, and $\\chi^{2}$=\\FortneyThreeChi \\ for $f$=$\\frac{1}{2}$ without TiO/VO). If the temperature inversion is due to TiO/VO, by the time the atmosphere becomes hot enough that TiO/VO remains in gaseous form, the thermal emission is too bright to fit the 3.6, and 5.8 $\\mu m$ thermal emission ($\\chi^{2}$=\\FortneyFourChi \\ for $f$=$\\frac{1}{2}$ with TiO/VO).  The combination of our blackbody and radiative transfer models with our own eclipse depth and those from the Spitzer/IRAC instrument \\citep{ODonovan09} thus suggest that the atmosphere of TrES-2b likely features modest redistribution of heat from the day to the nightside. It is unclear whether the atmosphere of TrES-2b requires a temperature inversion. A simple blackbody model ($f$=0.346 and $T_{eq}$$\\sim$1622 K) provides an exemplary fit to the data; this may indicate that TrES-2b has a fairly isothermal dayside temperature structure, perhaps similar to HAT-P-1b \\citep{Todorov10}. An important caveat, on the above result is that our $f$=$\\frac{1}{2}$ model without TiO/VO ($\\chi^{2}$=\\FortneyThreeChi) and thus without a temperature inversion returns nearly as good of fit as our best-fit blackbody model ($f$=0.346; $\\chi^{2}$=\\BlackbodyTwoChi); thermal emission measurements at other wavelengths, and repeated measurements at the above wavelengths, are thus necessary to differentiate a blackbody-like spectrum, from significant departures from blackbody-like behaviour, and to confirm that TrES-2b efficiently redistributes heat to the nightside of the planet.  Specifically, the variations between the models displayed in Figure \\ref{FigModel} are largest in the near-infrared J \\& H-bands and thus further near-infrared constraints -- if they are able to achieve sufficient accuracy to measure the small thermal emission signal of TrES-2b in the near-infrared -- should prove eminently useful to constrain the atmospheric characteristics of this planet.  If the excess emission at 4.5 $\\mu m$ is due to water emission, rather than absorption,  due to a temperature inversion in the atmosphere of TrES-2b then the inversion is unlikely to be due to TiO/VO. This is because the atmosphere of TrES-2b appears too cool to allow TiO/VO to remain in gaseous form in its upper atmosphere. If there is a temperature inversion then the high altitude optical absorber is likely to be due to another chemical species than TiO/VO. For instance, \\citet{Zahnle09} have investigated the photochemistry of sulphur-bearing species as another alternative.  TrES-2b is a promising target for the characterization of its thermal emission across a wide wavelength range. In addition to orbiting a relatively bright star, and having a favourable planet-to-star radius ratio, TrES-2 lies within the Kepler field. The combination of secondary eclipse measurements already published using Spitzer/IRAC, upcoming measurements with Kepler ($\\sim$430 - 900 $nm$; \\citealt{Borucki08}), and J, H and K-band near-infrared measurements that could be obtained from the ground, will allow us to fully constrain TrES-2b's energy budget. At the shorter end of this wavelength range it should also be possible to constrain the combination of reflected light and thermal emission. Our results predict that even if the geometric albedo of TrES-2b is as low as 5\\% in the Kepler bandpass, if Kepler is able to detect the secondary eclipse of this planet then it will be detecting a significant fraction of reflected light in addition to thermal emission. This will largely break the degeneracy on the Bond albedo and the reradiation factor for this planet, facilitating a more complete understanding of its energy budget."
    },
    "1005/1005.1091_arXiv.txt": {
        "abstract": "We, first, analytically work out the long-term, i.e. averaged over one orbital revolution, perturbations on the orbit of a test particle moving in a local Fermi frame induced therein by the cosmological tidal effects of the inhomogeneous Lema\\^{\\i}tre-Tolman-Bondi (LTB) model. The LTB solution has recently attracted attention, among other things, as a possible explanation of the observed cosmic acceleration without resorting to dark energy. Then, we phenomenologically constrain both the parameters $K_1\\doteq -\\ddot{\\mathfrak{R}}/{\\mathfrak{R}}$ and $K_2\\doteq -\\ddot{\\mathfrak{R}}^{'}/{\\mathfrak{R}}^{'}$  of the LTB metric in the Fermi frame by using different kinds of solar system data. The corrections $\\Delta\\dot\\varpi$ to the standard Newtonian/Einsteinian precessions of the perihelia of the inner planets recently estimated with the EPM ephemerides, compared to our predictions for them,  yield $K_1 = (4\\pm 8)\\times 10^{-26}$ s$^{-2}$, $K_2 = (3\\pm 7)\\times 10^{-23}$ s$^{-2}$. The residuals of the Cassini-based Earth-Saturn range, compared with the numerically integrated LTB range signature, allow to obtain  $K_1\\approx K_2\\approx 10^{-27}$ s$^{-2}$. The LTB-induced distortions of the orbit of a typical object of the Oort cloud with respect to the commonly accepted Newtonian picture, based on the observations of the comet showers from that remote region of the solar system, point towards  $K_1\\approx K_2\\lesssim 10^{-30}-10^{-32}$ s$^{-2}$. Such figures have to be compared with those inferred from cosmological data which are of the order of $K_1\\approx K_2=-4\\times 10^{-36}$ s$^{-2}$. ",
        "introduction": "Inhomogeneous cosmological models \\citep{Hellaby} have recently attracted much attention because, among other things, some of them may potentially be useful in explaining the observed Universe's acceleration without resorting to  dark energy. Thus, it is important to put them on the test independently of the phenomenon itself for which they have been purposely introduced.  In the framework of the general theory of relativity, the most general form of the line element $(ds)^2$ of a spherically symmetric inhomogeneous space-time in which the source in the Einstein's field equations is a perfect fluid is, in\\footnote{Recall that, by definition, the comoving coordinates are those in which the vector field $u^{\\alpha},\\ \\alpha=0,1,2,3$ has only the time component, i.e. $u^{\\alpha}\\propto \\delta^{\\alpha}_{\\ 0}$. If the vector field $u^{\\alpha}$ has zero rotation, as it happens if the space-time is spherically symmetric and the metric obeys the Einstein equations with a perfect fluid source, the comoving coordinates can be chosen so that they are synchronous, i.e. in them the metric tensor has no off-diagonal components $g_{0i},\\ i=1,2,3$.} comoving-synchronous spherical coordinates $\\{t,r,\\theta,\\phi\\}$ \\citep{Krasi}, {{\\eqi (ds)^2 = e^{{\\mathfrak{C}}(r,t)}(cdt)^2-e^{\\mathfrak{A}(r,t)}(dr)^2 - {\\mathfrak{R}}^2(r,t)[(d\\theta)^2+\\sin^2\\theta (d\\phi)^2].\\lb{metrica}\\eqf}} In order to solve the resulting Einstein field equations, an equation of state has to be assumed. The most natural choice consists of setting the pressure equal to zero, corresponding to a dust evolution driven by gravitation only. From the equations of motion of a perfect fluid it turns out that, for $p=0$, the dust moves along timelike geodesics. As a consequence, $\\mathfrak{C}=0$ and, for\\footnote{Here and in the following,  $\\mathfrak{R}^{'}$ denotes the partial derivative of $\\mathfrak{R}$ with respect to $r$.} ${\\mathfrak{R}}^{'}\\neq 0$, \\eqi e^{\\mathfrak{A}}=\\rp{   {{\\mathfrak{R}}^{'}}^2  }{1+2\\mathfrak{E}(r)},\\eqf where $\\mathfrak{E}(r)$ is an arbitrary function such that $\\mathfrak{E}\\geq -1/2$ for all $r$ in order to preserve the right signature of the metric of \\rfr{metrica}. The resulting space-time line element is the so-called Lema\\^{\\i}tre-Tolman-Bondi (LTB) model. Indeed, such spherically symmetric inhomogeneous dust models were first discovered by \\citet{Lem} and further studied by \\citet{Tol} and \\citet{Bon}; for a detailed discussion of several properties see \\citet{Krasi}. Such models are among the best known and most useful exact solutions of Einstein's equations. Since they allow us to examine non-linear effects analytically, or at least in a tractable way, there is an extensive literature using them mostly to describe cosmological inhomogeneities \\citep{Hella}, but also as in other theoretical contexts, such as gravitational collapse \\citep{colla} and censorship of singularities \\citep{Cens} or quantum gravity \\citep{quantum}. Several authors considered LTB models as tools to probe how the cosmic acceleration associated to recent observations can be accounted for inhomogeneities, without introducing dark energy \\citep{Paran,Sark,Enq,Bellido}. LTB models are also a standard choice  to apply Buchert's scalar averaging formalism \\citep{avera}, in which the effects of dark energy could be mimicked by \\virg{backreaction} terms. See \\citet{Noelle} for a comprehensive review of such aspects.  Here we are interested in putting constraints in a phenomenological way on some parameters of the LTB metric from its local tidal effects on the dynamics of test bodies of the solar system.  To this aim, it is necessary to set up a standard orthonormal tetrad frame parallel transported along the worldline of a fundamental observer representing, in this case, the Sun, and explicitly derive the expression of the LTB metric in such a local Fermi frame. This has been recently done by \\citet{Mash07}. In general, the motion of test bodies can be studied in such  Fermi coordinate systems following ideas and methods developed by \\citet{Syn,Mana,Mash}.  The paper is organized as follows. In Section \\ref{dua} we analytically work out the LTB effects on the orbital motion of a test particle in the local Fermi frame and compare them with some results present in literature. In Section \\ref{tria} we compare our theoretical predictions to different types of bodies and data of the solar system to phenomenologically put constraints on \\citep{Mash07} \\begin{eqnarray} K_1 &\\doteq & -\\rp{\\ddot{\\mathfrak{R}}}{\\mathfrak{R}}, \\\\ \\nonumber \\\\ K_2 &\\doteq & -\\rp{\\ddot{\\mathfrak{R}}^{'}}{\\mathfrak{R}^{'}} \\end{eqnarray} entering the LTB metric written in terms of the Fermi coordinates of a fundamental observer. Concerning their values infered from cosmological data,  from (62)-(64) by \\citet{Mash07} it turns out that, in proximity of the origin of the local Fermi frame, i.e. for $r\\rightarrow 0$, $K_1$ and $K_2$ are rather similar being of the order of\\footnote{$K_{1/2}$ means that both $K_1$ and $K_2$ have approximately the same values.} \\eqi K_{1/2}\\approx q_0 H_0^2.\\eqf In it \\citep{hubblo} \\eqi H_0=71.0 \\ {\\rm km s}^{-1}\\ {\\rm Mpc}^{-1}=2.3\\times 10^{-18}\\ {\\rm s}^{-1}\\eqf is the present value of the Hubble parameter, and \\citep{Xu} \\eqi q_0 \\doteq\\rp{1+3\\widetilde{\\Omega}_{\\rm DE}w_{\\rm DE}}{2}=-0.7\\eqf is the deceleration parameter in which \\citep{hubblo} $\\widetilde{\\Omega}_{\\rm DE}=0.734\\pm 0.029$ is the dark energy density and \\citep{hubblo} $w_{\\rm DE}=-1.12^{+0.42}_{-0.43}$ is the equation of state. Thus, the order of magnitude of $K_1$ and $K_2$ inferred from cosmological observations is \\eqi K_{1/2}\\approx -4\\times 10^{-36}\\ {\\rm s}^{-2}.\\eqf Section \\ref{quatra} is devoted to summarizing our findings and to the conclusions. ",
        "conclusions": "\\lb{quatra} We, first, analytically worked out the long-term, i.e. averaged over one orbital revolution, perturbations on the orbit of a  test particle moving in a local Fermi frame induced by the cosmological tidal effects of the inhomogeneous Lema\\^{\\i}tre-Tolman-Bondi model which has recently attracted attention as possible explanation of the observed cosmic acceleration without resorting to dark energy. In particular, we computed the variations of the semimajor axis $a$, the eccentricity $e$, the inclination $I$, the longitude of the ascending node $\\Omega$, the argumet of pericenter $\\omega$ and the mean anomaly $\\mathcal{M}$ by means of the Lagrange's variational equations with the eccentric anomaly $E$ as fast variable. We found that $a$ does not experience any change, on average, while the other Keplerian orbital elements undergo long-term variations including both secular and harmonic terms with frequency $2\\omega$; $e$ and $I$ exhibit only sinusoidal changes, while $I,\\Omega,\\mathcal{M}$  secularly precess as well. We repeated the calculation also with the Gauss perturbative approach  by finding the same results. We also explicitly computed the changes  over one orbital revolution of the position and velocity vectors along the radial $\\bds{\\hat R}$, transverse $\\bds{\\hat \\Xi}$ and normal directions $\\bds{\\hat \\Upsilon}$ of a frame co-moving with the test particle. While the radial and normal components  $\\Delta R$ and $\\Delta\\Upsilon$ of the perturbation of the position vector experience only harmonic variations, the transverse one $\\Delta\\Xi$ shows secular changes as well. Concerning the velocity, its radial perturbation $\\Delta  V_R$ undergoes both secular and sinusoidal modifications, while its transverse and normal components $\\Delta V_{\\Xi}$ and $\\Delta V_{\\Upsilon}$ exhibit harmonic signatures only. In obtaining our results we  made no approximations on $e$ and $I$ by retaining all terms.  Then, we phenomenologically put constraints on both the parameters $K_1$ and $K_2$ of the LTB metric in the local Fermi frame  by looking at various astronomical bodies and data of the solar system. By comparing our analytical prediction for the rate of the longitude of pericenter $\\varpi$ to different sets of the corrections $\\Delta\\dot\\varpi$ to the standard Newtonian/Einsteinian precessions of the perihelia of the inner planets recently estimated with the EPM ephemerides we found that the tightest constraints are $K_1 = (4\\pm 8)\\times 10^{-26}$ s$^{-2}$, $K_2 = (3\\pm 7)\\times 10^{-23}$ s$^{-2}$. The confrontation of the residuals of the Earth-Saturn range, obtained by processing some years of radiotechnical data from the Cassini spacecraft as well, with the numerically computed LTB-induced Earth-Saturn range signal allowed to set  $K_1\\approx K_2\\approx 10^{-27}$ s$^{-2}$. By looking at the LTB-type distortions of the orbit of a typical object of the Oort cloud with respect to the commonly accepted Newtonian picture, based on the observations of the comet showers from that remote region of the solar system, pointed towards even lower bounds, i.e. $K_1\\approx K_2\\lesssim 10^{-30}-10^{-32}$ s$^{-2}$. According to cosmological data, $K_1\\approx K_2=-4\\times 10^{-36}$ s$^{-2}$."
    },
    "1005/1005.4999_arXiv.txt": {
        "abstract": "A time-dependent Synchrotron Self Compton model (SSC) which is able to motivate the used electron spectra of many SSC models as a balance of acceleration and radiative losses is introduced. Using stochastic acceleration as well as Fermi-I processes even electron spectra with a rising part can be explained, which are mandatory to fit the lowstate spectral energy distribution (SED) of PKS 2155-304 as constrained from Fermi LAT observations. Due to the time resolution the outburst of PKS 2155-304 observed by H.E.S.S. in 2006 can be modelled selfconsistently as fluctuations along the jet axis without introducing new sets of parameters. The model makes the time evolution of the SED also accessible. Hence giving new insights into the flaring behavior of blazars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.1572_arXiv.txt": {
        "abstract": "Using three epochs of VLA observations of the Galactic Plane in the first quadrant taken $\\sim15$ years apart, we have conducted a search for a population of variable Galactic radio emitters in the flux density range 1--100 mJy at 6~cm. We find 39 variable sources in a total survey area of 23.2 deg$^2$. Correcting for various selection effects and for the extragalactic variable population of active galactic nuclei, we conclude there are $\\sim1.6$ Galactic sources deg$^{-2}$ which vary by more than 50\\% on a time scale of years (or shorter). We show that these sources are much more highly variable than extragalactic objects; more than 50\\% show variability by a factor $>2$ compared to $<10\\%$ for extragalactic objects in the same flux density range. We also show that the fraction of variable sources increases toward the Galactic center (another indication that this is a Galactic population), and that  the spectral indices of many of these sources are flat or inverted. A small number of the variables are coincident with mid-IR sources and two are coincident with X-ray emitters, but most have no known counterparts at other wavelengths. Intriguingly, one lies at the center of a supernova remnant, while another appears to be a very compact planetary nebula; several are likely to represent activity associated with star formation regions. We discuss the possible source classes which could contribute to the variable cohort and followup observations which could clarify the nature of these sources. ",
        "introduction": "Variable radio emission is a hallmark of energetic objects such as coronally active stars, supernovae, neutron stars, black holes, and active galactic nuclei (AGN). Indeed, radio variability is often indicative of high-energy processes and, in principle, can be valuable for finding examples of relatively rare objects. However, surveys for variability are themselves quite rare --- blind sky surveys are almost never repeated owing to the scarcity of telescope time. The exceptions are mostly in the optical regime. Comparisons between POSS1, POSS2, and SDSS have been useful for studying variability (de~Vries et al.\\ 2005). Gravitational microlensing studies (e.g., Alcock et al.\\ 1997) and supernova searches (e.g., Astier et al.\\ 2006, Miknaitis et al.\\ 2007) have produced a wealth of data on optical variability from targeted sky regions, and several upcoming experiments such as Pan-STARRS (Kaiser et al.\\ 2002), the Palomar Transit Factory (Rau et al.\\ 2009), and LSST (Tyson 2002) will make the coming decade one in which time-domain astronomy plays a prominent role.  Variability studies in the radio band have typically targeted bright extragalactic sources (see de~Vries et al.\\ 2004 for a review of searches for, and mechanisms of, radio variability). Comparisons between blind radio surveys are often hampered by differences in angular resolution and the confusing presence of interferometric sidelobe patterns. For example, there has been no systematic search for radio variability between the two largest radio sky surveys, FIRST (Becker et al.\\ 1995) and NVSS (Condon et al.\\ 1998). The FIRST survey did observe one area twice at 1400 MHz, an equatorial strip $\\sim1.5$ degrees wide in the range $21^{\\mathrm h}20^{\\mathrm m} < RA < 03^{\\mathrm h} 20^{\\mathrm m}$. A search for variable sources in this area was reported in de~Vries et al.\\ (2004). The search covered $\\sim120$~deg$^2$ of extragalactic sky with a sensitivity similar to the Galactic plane search reported here; it thus serves as a useful control from which to estimate how many of the sources we find are background extragalactic radio sources.  The most systematic search for radio variablity in the Galactic plane used the NRAO 91-m telescope in Greenbank, WV, operating at a frequency of 5~GHz (Gregory \\& Taylor 1986). Over a five-year period the plane was observed 16 times, leading to the detection of 32 variable radio sources. The survey had a flux density threshold of $\\sim20$~mJy and an angular resolution of 3\\arcmin. Using the Very Large Array\\footnote{The Very Large Array is an instrument of the National Radio Astronomy Observatory, a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} (VLA), the Galactic plane has been surveyed at this same frequency but with much higher sensitivity and angular resolution, although with minimal repetition (Becker et al.\\ 1994).  Recently, a new Galactic plane survey at 6~cm (4.86~GHz) has begun at the VLA. The new survey (CORNISH\\footnote{The Co-Ordinated Radio 'N' Infrared Survey for High-mass star formation; see \\url{http://www.ast.leeds.ac.uk/Cornish}}; Purcell et al.\\ 2008) has substantial overlap with our previous survey; both surveys have a flux density threshold of $\\sim1$~mJy. Comparison between these two data sets allows for a search for Galactic radio sources that exhibit variability over the fifteen-year interval between the surveys. The search is complicated because the two surveys use different VLA configurations and hence have different angular resolutions (5\\arcsec\\ versus 1.5\\arcsec). Nonetheless, it is possible to identify strongly varying sources. In this paper we will compare results between the original survey and two epochs of data from the new survey. In section 2 we discuss the parameters of the two surveys, while in section 3 we present the results from a comparison of the two samples and adduce evidence for variability. We describe the properties of the variable sources including their spatial distribution, spectral indices, and counterparts at other wavelengths (\\S4) and end with a discussion of our limited knowledge of the nature of these objects (\\S5).  \\begin{figure*} \\epsscale{1.0} \\plotone{f1.eps} \\caption{ Sky coverage for overlapping regions from the 1990${+}$ (red), 2005 (blue), and 2006 (gray) 6~cm survey epochs. Darker regions have higher rms noise values, while white areas are outside the survey. The noise is higher at the edges of the coverage and in fields with bright or complex extended sources.  Typical rms values are $\\sim 0.1$~mJy in both surveys, but the old 6~cm data were acquired with a more widely spaced pointing grid and so display greater variation with position. The areas of overlap are given in Table~1. } \\label{fig-coverage} \\end{figure*} ",
        "conclusions": "\\subsection{Source Identification -- What They Are Not}  Having established the existence of a population of highly variable Galactic sources, the obvious question is, What are they? Three classes of Galactic variable radio emitters can be easily eliminated from consideration: coronally active radio-emitting stars, pulsars, and masers. We justify our exclusion of these source classes in turn.  In a survey of 122 RS CVn and related active binary systems which are among the most luminous stellar radio sources, Drake et al. (1989) found only 18 detected above a quiescent flux density of 1~mJy at 6~cm; the faintest optical counterpart was $V=10.0$. Even assuming an extreme flare of a factor of 100 (Osten 2008), the faintest possible counterpart would have $V=15$; none of our variables has a counterpart this bright. As for dMe flare stars, the other main class of variable stellar radio emitters, the most luminous quiescent emission is $\\sim10^{14.2}$ erg s$^{-1}$ Hz$^{-1}$ (Gudel et al.\\ 1993) corresponding to a flux density of $\\sim1$~mJy at a distance of 13~pc. Even an extreme flare with an increase of a factor of 500 over the quiescent level (Osten 2008 and references therein) would fall below our flux density threshold for a distance $>290$~pc. Stars with spectral types later than M6 would have counterparts fainter than 20th magnitude and could be represented in our sample. However, statistically, M-stars cannot be a significant contributor; Helfand et al.\\ (1999) found only $\\sim5$ M stars in 5000 deg$^2$ of the {\\sl FIRST} survey to a flux density limit of 0.7~mJy, whereas our variables have a surface density of 1.6~deg$^{-2}$.  While nearby radio pulsars scintillate strongly in the ISM leading to large-amplitude variability, pulsars have very steep radio spectra, and most have not been detected at 6~cm (none of our objects are coincident with one of the 1827 known pulsars; Manchester et al. 2005). For a typical spectral index of $-1.5$, our weakest source would be a $\\sim100$~mJy pulsar at 400~MHz, and most unlikely to have been missed in pulsar surveys.  The small duty cycle of the recently discovered RRATs (Rotating RAdio Transients -- McLaughlin et al.\\ 2006) makes them equally unlikely to explain our variable sources.  Finally, radio masers are known to be highly variable, but no known maser transitions fall within our bandpass. As noted below, however, three of our variables are coincident with methanol masers.  Two classes of extragalactic radio transients --- supernovae and GRB afterglows --- are also highly improbable counterparts for our events. Both have rise times of at most tens of days and cannot, in the absence of a steady underlying source of radio emission, account for the bulk of our sources which show a flux density increase over many years. In addition, their rarity makes them statistically unlikely counterparts. The one extragalactic population that does show variability on the time scales we probe, AGN, are shown above to have variability amplitudes which exclude them from explaining all but a handful of our events.  The remaining known classes of variable radio sources include microquasars (accreting, high-mass X-ray binaries that produce relativistic jets: e.g., SS433, Cyg X-3 and GRS 1915${+}$105), radio magnetars (Camilo et al.\\ 2006; Camilo et al.\\ 2007), and the recently described Galactic Center Transient sources (Hyman et al.\\ 2009 and references therein).  The first two of these have signatures at other wavelengths; we explore below the fragmentary data outside the radio band that is available for our variable objects.  \\subsection{Source Identification -- Multiwavelength Data}  Counterparts at other wavelengths can be useful in suggesting the origin of radio variability. At our MAGPIS website (Helfand et al. 2006), we have collected the following Galactic plane data in addition to the three-epoch 6~cm data described herein: two epochs of 20~cm observations for these same fields including the principal MAGPIS survey, 90~cm observations of the same regions, the 3.6, 4.5, 5.8, and 8.0 $\\mu$m data from the Spitzer {\\sl GLIMPSE} survey (Benjamin et al.\\ 2003), 24 $\\mu$m images from MIPSGAL (Carey et al.\\ 2009), 20 $\\mu$m data from the MSX survey (Price et al.\\ 2001) and the 1.1~mm Bolocam Galactic plane survey (Aguirre et al.\\ 2009). In addition, we have queried the SIMBAD database for each of our sources and have examined the Digitized Sky Survey images; in one case, we have obtained optical observations of a source. We report the results of this multi-wavelength inquiry here.  \\subsubsection{Mid-IR and mm observations}  Seven of our variables are detected at 24 $\\mu$m in the MIPSGAL survey, and six of these are also detected in at least one GLIMPSE mid-IR band. Four of the objects are also detected at 1.1~mm in the Bolocam survey. In all seven cases at least two bands are available, and in all seven cases the sources are red; i.e., they are faintest in the short-wavelength bands and brightest in the long-wavelength bands. In two cases (G31.1595${+}$0.0449 and G37.7347${-}$0.1126) multiple components with different IR spectral shapes are present, with the radio source identified with the brighter component in the first case, and the redder component in the second. Three of the IR-detected objects have associated methanol masers; this, coupled with their IR spectra demonstrate they represent activity associated with star formation in compact or ultracompact \\ion{H}{2} regions.  For one IR-detected source, G29.5779${-}$0.2685, we have obtained followup observations at the MDM Observatory (J.~Halpern, private communication).  R-band and H$\\alpha$ images were obtained on 23~August 2009, and show a barely resolved ($\\sim1\\arcsec$) object, brighter in H$\\alpha$ and coincident with the radio source. A spectrum obtained the same night shows no continuum, but very strong nebular emission lines. The object appears to be a very compact planetary nebula. Its radio flux history is thus perplexing: 6.9~mJy at 6~cm in $\\sim1990$, rising to 10.5~mJy in 2005 and falling again to 5.8~mJy in 2006. The 20~cm flux density in the MAGPIS survey (epoch 2001--2004) is only 1.3~mJy, suggesting the source may be optically thick. Further simultaneous multi-frequency observations are required to measure the radio spectrum and derive clues as to the nature of the source's variability.  \\begin{figure*} \\epsscale{1.0} \\plotone{f7.eps} \\caption{ MAGPIS 20~cm image of supernova remnants W41 (G23.3${-}$0.3) and G22.7${-}$0.2.  The boxes mark the positions of three variable 6~cm sources (G22.7194${-}$0.1939, G22.9116${-}$0.2878, and G22.9743${-}$0.3920). } \\label{fig-snrs} \\end{figure*}  \\subsubsection{X-ray observations}  The brightest variable, G21.6552${-}$0.3611, is coincident with a point-like X-ray source catalogued in the XMM Galactic Plane survey (Hands et al.\\ 2004). It has a hard band (2--6~keV) flux of 0.0051 ct s$^{-1}$ and is undetected in the soft (0.4--2.0~keV) band. For an intrinsic power-law spectrum with spectral index $\\Gamma = 1.9$, the expected absorption column density through the Galactic plane of $\\sim10^{23}$ cm$^{-2}$ is consistent with the non-detection in the soft band; the inferred intrinsic flux in the 0.2--10~keV band would be $7 \\times 10^{-13}$ erg cm$^{-2}$ s$^{-1}$. For an extragalactic AGN at 1 Gpc, this corresponds to a luminosity of $8 \\times 10^{42}$ erg s$^{-1}$, while for a Galactic object at 5~kpc, the X-ray luminosity would be a modest $2 \\times 10^{33}$ erg s$^{-1}$; for a column density of only $10^{22}$ cm$^{-2}$, the luminosity estimates are lower by a factor of 3. While this source is the brightest of our variables, it has one of the lowest modulation factors (decreasing by just $\\sim50\\%$ over 16 years). The (non-contemporaneous) 20~cm flux density is lower than either of the 6~cm values, suggesting a mildly inverted spectrum source. It is not detected at any other wavelength. The most likely explanation of this object is a flat-spectrum extragalactic radio source, one of a handful we expect in our sample.  One other source, G30.4460${-}$0.2148, lies 27\\arcsec\\ from the position of an ASCA Galactic Plane Survey catalog entry (Sugizaki et al.\\ 2001). The uncertainty in the X-ray position is 1\\arcmin; one other (brighter) radio source lies within the X-ray error circle although at twice the distance from its centroid. The X-ray source is a marginal detection ($4.6 \\sigma$) with a 0.7--7.0~keV unabsorbed flux of $2.6 \\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$ for an intrinsic power law index of $\\Gamma = 1.9$ and an absorption column density of $10^{23}$ cm$^{-2}$; the flux is roughly four times lower for $N_H = 10^{22}$ cm$^{-2}$. Assuming the identification is correct, the X-ray to radio flux ratio is thus $\\sim20$ times greater than our other X-ray detection, although still within the X-ray to radio luminosity ratios characteristic of AGN. The primary distinguishing feature, however, is that the radio source is coincident with a very bright mid-IR source (saturated in all but the 3.6 $\\mu$m band) which is also detected at 1.1~mm.  The ASCA Galactic Plane Survey covered an area encompassing all but six of our variables to a flux density level of approximately $10^{-12.5}$ erg cm$^{-2}$ s$^{-1}$; no other X-ray sources are coincident to within 1\\arcmin. The higher-resolution coverage of the Einstein, ROSAT, XMM, and Chandra is much spottier; no further matches are found within the 10\\arcsec\\ error circles of these other catalogs.  \\subsubsection{Low-frequency radio detections}  Three of the variable sources are detected at 90~cm. G22.9116${-}$0.2878 (Fig.~\\ref{fig-snrs}) has a 90~cm flux density of $\\sim180$~mJy; this is consistent with a nonthermal spectral index of $\\sim-0.9$ if one takes the most recent (but far from contemporaneous) 6, 20, and 90~cm measurements. The 6~cm flux density increased by more than a factor of three since 1990, making it one of the higher amplitude variables, but no other information is available on this source. G30.6724${+}$0.9637 (the highest latitude source detected) has a 90~cm flux density of $\\sim70$ mJy, below that of the 20~cm flux density (90~mJy), possible additional evidence for variability, as the 20:6~cm flux density ratio is 3:1 (again, all non-contemporaneous). This is the smallest amplitude variable in our sample and, given its distance from the Galactic plane, an extragalactic counterpart is the most likely explanation.  The third 90~cm detection is G22.7194${-}$0.1939, perhaps the most intriguing source in our sample. An image of the region surrounding this source is given in Figure~\\ref{fig-snrs}. The source lies very near to the geometric center of a 30\\arcmin-diameter supernova remnant, G22.7${-}$0.2 (Green 2004 and references therein) and 4\\arcmin\\ from a fairly bright \\ion{H}{2} region. There is no counterpart detected at mm, IR, or optical wavelengths. The source brightened by a factor of four between 2003 and 2006 at 6~cm; its 20~cm flux density is 12~mJy, three times higher than the higher of the two 6~cm measurements. The distance to the remnant is unknown, although its large angular diameter would suggest it is not very remote (its diameter would be $\\sim45$~pc at 5~kpc). X-ray observations could reveal whether or not this source is likely to be a compact object associated with the supernova remnant.  \\subsection{Summary}  We have discovered a relatively high surface density (2 deg$^{-2}$) of variable radio sources in the Galactic plane and have argued that the large majority of these ($\\sim80$\\%) are Galactic objects. While a few are associated with young star formation activity, the identity of the majority is unknown. Follow up radio observations are required to confirm the variability in these sources, establish the variability time scale(s), and obtain contemporaneous spectral indices. Observations at optical, infrared, and X-ray wavelengths could help establish counterparts and identify the origin of the variable radio emission."
    },
    "1005/1005.4161_arXiv.txt": {
        "abstract": "The \\r-mode instability in rotating compact stars is used to constrain the phase of matter at high density. The color-flavor-locked phase with kaon condensation (CFL-K0) and without (CFL) is considered in the temperature range $10^8$K $\\lesssim T\\lesssim 10^{11}$K. While the bulk viscosity in either phase is only effective at damping the \\r-mode at temperatures $T\\gtrsim10^{11}$K, the shear viscosity in the CFL-K0 phase is the only effective damping agent  all the way down to temperatures $T\\gtrsim 10^8$K characteristic of cooling neutron stars. However, it  cannot keep the star from becoming unstable to gravitational wave emission for rotation frequencies $\\nu\\approx 56-11$Hz at $T\\approx10^8-10^9$K. Stars composed almost entirely of CFL or CFL-K0 matter are ruled out by observation of rapidly rotating neutron stars, indicating that dissipation at the quark-hadron interface or nuclear crust interface must play a key role in damping the instability. ",
        "introduction": "Introduction}   The $r$-mode instability of rotating compact stars is of wide astrophysical interest for several reasons~\\cite{Andersson:1997xt,*Friedman:1997uh,*Lindblom:1998wf}. The majority of neutron stars have rotation period $P\\sim 0.5$seconds, much slower than theoretical limits: $P\\sim 0.001$s ($\\nu=1/P\\sim 1$kHz)~\\cite{Shapiro:1984}. Even the fastest spinning neutron star PSR J1748-2446ad with $\\nu=716$Hz~\\cite{Hessels:2006} is not near the theoretical stability limit, with the majority of  ``spun-up'' neutron stars in low mass X-ray binaries (LMXBs) catalogued between 300-640 Hz. The \\r-mode provides a mechanism for spinning down young neutron stars~\\cite{Andersson:1998qs} and limiting spin-up frequencies in millisecond pulsars (see eg.~\\cite{Lorimer:2008se} for a review) through angular momentum loss in gravitational wave emissions. These gravitational waves may be detected given the expected improved sensitivities of ground-based interferometers eg. VIRGO, advanced LIGO~\\cite{Sa:2007zz,*Watts:2008qw,*Stergioulas:lrr-2003-3}. In addition, \\r-modes are probes of the equation of state of cold and dense matter  through viscous damping effects and advance our understanding of nuclear and quark interactions by providing an astrophysical and phenomenological link to Quantum Chromodynamics (QCD) over a certain range of temperatures ($10^7$K $\\lesssim T\\lesssim10^{11}$K) and at super-saturation densities [$n_B\\sim$ (2-5)$n_{\\rm sat}$]. Only a few studies of \\r-modes linking them to phenomenology of quark matter in neutron stars have been carried out thus far~\\cite{Madsen:1999ci,Andersson:2001ev,Jaikumar:2008kh,Sa'd:2008gf,Mannarelli:2008je}. In this work we use \\r-mode damping to constrain the presence of a realistic phase of superconducting quark matter~\\cite{Alford:1998mk,Rapp:1997zu} in the core of neutron stars: viz., one that includes the effects of a sizable strange quark mass $m_s$ at intermediate baryon density.  A key parameter in understanding the \\r-mode instability in rotating compact stars is the critical frequency $\\Omega_c$ for the onset of the instability ($\\Omega=2\\pi/P$~rad/s).  Gravitational radiation tends to make the $r$-mode grow on a characteristic time scale $\\tau_{GR}$. Internal friction such as viscosity or mutual friction tends to damp $r$-mode growth on a characteristic time scale $\\tau_F$. The competition between the two determines $\\Omega_c$ at any temperature. At $\\Omega\\geq\\Omega_c$, the $r$-mode develops an instability ($1/\\tau_{\\rm total}=1/\\tau_{GR}+1/\\tau_{F}<0$). Previous works~\\cite{Sa'd:2008gf,Jaikumar:2008kh,Mannarelli:2008je} have studied the \\r-mode damping in normal (ungapped) quark matter as well as some superconducting (gapped) phases such as the color-flavor-locked phase (CFL), and found that bulk and shear viscosity in the  fully gapped phase is large enough to render rapidly spinning quark stars ($\\Omega\\geq 0.1\\Omega_\\mathrm{Kepler}$) stable.  However, this conclusion applies only in a restricted temperature regime $T\\gtrsim 10^{10}$K, since the shearing mean free path due to phonons otherwise becomes larger than the star's radius ($\\sim 10$km). Furthermore, for temperatures $T\\lesssim 10^8$K, mutual friction in the CFL phase has been shown to be too weak to damp \\r-modes~\\cite{Mannarelli:2008je}, arguing for rapid spin-down of CFL stars to less than a few Hz, effectively ruling our cold CFL matter in any rapidly rotating neutron stars. The \\r-mode damping in the temperature regime $10^8$K $\\lesssim T \\lesssim10^{10}$K for color-superconducting phases of quark matter is therefore an open question, and is the focus of this work. Here we study $r$-mode damping and the critical frequency curve for a neutron star made up mostly or entirely of the kaon condensed CFL (CFL-K0) phase.  The CFL phase with symmetric pairing of up-, down- and strange-quarks is severely stressed at realistic chemical potential $\\mu_q\\sim 300$ MeV because the strange quark mass $m_s\\sim 100$ MeV is non-negligible and costs extra energy $m_s^2/(2\\mu_q)$ compared  to the much lighter up- and down-quarks. In Ref.~\\cite{Bedaque:2001je}, it was shown that it is energetically favorable for the CFL vacuum to be in a rotated chiral state with a non-zero kaon condensate when the cost $m_s^2/(2\\mu_q)$ exceeds the kaon mass $m_K$, the lightest meson in high-density QCD~\\cite{Son:1999cm}. Realistic neutron star densities could support the CFL-K0 phase. We find that in the CFL-K0 phase, while the bulk viscosity has little role to play in the temperature region relevant to \\r-mode damping in cooling neutron stars, the shear viscosity from kaons in the CFL-K0 phase is important even below $T\\lesssim 10^9$K, unlike the CFL phase. In essence, this fact coupled with the large mean free path of the phonons controls the main features of the critical frequency curve (Fig.~\\ref{fig:critfreq}). ",
        "conclusions": "Conclusions}  \\r-mode oscillations of compact stars in the kaon-condensed CFL phase are considered. The mode frequency is almost exactly the same as that for a CFL star~\\cite{Jaikumar:2008kh} since the equation of state for the CFL-K0 phase only differs at ${\\cal O}(m_S^4)$.  The mode growth and viscous damping timescales in the temperature range  $10^8$K $\\lesssim T\\lesssim 10^{11}$K is studied based on the dominant bulk and shear viscosity contributions. Compared to the pure CFL phase, the bulk viscosity in CFL-K0 is smaller and does not play a significant role. At temperatures $T\\gtrsim 10^{10}$K shear viscosity associated with phonon scattering is important and determines the temperature dependence of the critical frequency curve. At lower temperatures, the phonon mean free path exceeds the star radius, so that the new mechanism for shear viscosity associated with kaon condensation~\\cite{Alford:2009jm} provides the dominant dissipation. In comparison to pure CFL stars, CFL-K0 stars are more stable against the \\r-mode instability at the lower end of the temperature range considered. However, given the observed LMXB spin rates, the critical frequency curve predicts that even pure CFL-K0 stars are unlikely to exist. Just as in the neutron matter case, viscous damping just beneath the core-crust interface (Ekman layer) appears to be necessary to provide a large enough stable rotation rate consistent with LMXB data~\\cite{Bildsten:1999zn,*Levin:2000vq}.  The main result of the current work is that given the known dissipative mechanism in CFL or CFL-K0 phase, a pure quark star with these phases is ruled out by observed LMXB spin rates. Mutual friction associated with kaon-vortex scattering could provide an additional dissipative mechanism, but the critical frequency curves calculated here provide a lower bound. Future related work should investigate hybrid stars with a CFL-K0 core. The role of the crust-core interface as well as the quark-hadronic matter layer should be key ingredients in obtaining constraints on the persistence of gravitational waves from stars containing superfluid quark matter in their interior."
    },
    "1005/1005.1791.txt": {
        "abstract": "We present a new set of cooling models and isochrones for both H- and He-atmosphere white dwarfs, incorporating accurate boundary conditions from detailed model atmosphere calculations, and carbon-oxygen chemical abundance profiles based on updated stellar evolution calculations from the BaSTI stellar evolution archive - a theoretical data center for the Virtual Observatory. We discuss and quantify the uncertainties in the cooling times predicted by the models, arising from the treatment of mixing during the central H- and He-burning phases, number of thermal pulses experienced by the progenitors, progenitor metallicity and the $^{12}C(\\alpha,\\gamma)^{16}O$ reaction rate. The largest sources of uncertainty turn out to be related to the treatment of convection during the last stages of the progenitor central He-burning phase, and the $^{12}C(\\alpha,\\gamma)^{16}O$ reaction rate.  We compare our new models to previous calculations performed with the same stellar evolution code, and discuss their application to the estimate of the age of the solar neighborhood, and the interpretation of the observed number ratios between H- and He-atmosphere white dwarfs. The new white dwarf sequences and an extensive set of white dwarf isochrones that cover a large range of ages and progenitor metallicities are made publicly available at the official BaSTI website. ",
        "introduction": "\\label{intro} % %%%%%%%%%%%%%%%%%%%%%%  The interpretation of photometric and spectroscopic observations of stellar populations relies on the use of grids of stellar models and isochrones, that have to cover a wide range of initial chemical compositions, stellar masses and evolutionary phases. The BaSTI (a Bag of Stellar Tracks and Isochrones) project\\footnote{Official website at \\url{http://www.oa-teramo.inaf.it/BASTI}} started in 2004 has delivered, to date, an homogeneous database of stellar evolution models, isochrones and integrated spectra for single-age, single-metallicity populations, encompassing a large chemical composition range appropriate for stellar populations harboured in star clusters and galaxies of various morphological types (Pietrinferni et al.~2004, 2006, 2009, Cordier et al.~2007, Percival et al.~2009). Results from BaSTI projects have been used by a large number of authors to address very diverse astrophysical problems like, among others, fitting eclipsing binary systems in the mass-radius plane, determining the ages of star clusters from their color-magnitude-diagrams (CMDs) or comparing integrated colors of elliptical galaxies with theoretical predictions.  BaSTI models and isochrones in their present form cover all relevant evolutionary phases until either the end of the thermal pulse regime along the Asymptotic Giant Branch (AGB), or central carbon ignition for masses without electron degenerate carbon-oxygen (CO) cores. In this paper we extend the evolutionary phase coverage of our database to include cooling models of CO-core White Dwarfs (WDs), the final evolutionary phase of stars with initial masses smaller than about 6-7 $M_{\\odot}$. %Typical CO-core WDs (hereafter denoted simply as WDs) have masses between $\\sim$0.55 and %$\\sim 1.0~M_{\\odot}$ and are made of a nearly %isothermal electron degenerate core surrounded by a layer of pure He with mass of the order %of $M_{He} \\sim 10^{-2} \\ M_{WD}$ (where $M_{WD}$ is the total mass of the WD) %or less. This He-layer can be, in turn , surrounded by a H-layer with mass %of the order of $M_H \\sim 10^{-4} \\ M_{WD}$ or less. %In absence of metal accretion from the environment, given the high efficiency of atomic diffusion %at the very beginning of the cooling sequence (see.e.g., Koester~2009 for %a tabulation of metal diffusion timescales in WD envelopes, as a function of $T_{eff}$) all metals in the %progenitor's envelope have settled at the bottom of the He-layer. %The WD evolution is a cooling process, whereby the core acts as energy reservoir (stored as %the internal energy of the CO ions) while the outer non-degenerate layers control the rate of energy outflow. As a result, %the luminosity and core temperature decrease with time, at approximately constant radius. %Given that most stars are or will become WDs, plus the existence of %a well defined relationship between cooling time and luminosity, and the large cooling timescales, WDs have been since long %recognized as very attractive candidates to unveil the history of star formation in the Galaxy (e.g., Schmidt~1959).  During the last two decades observations and theory have improved to a level that has made finally possible to employ WDs for determining ages of the stellar populations in the solar neighborhood (e.g., Winget et al.~1987, Garcia-Berro et al.~1988, Wood 1992, Oswalt et al.~1996), and in the nearest open (e.g., Richer et al.~1998, von Hippel~2005, Bedin et al.~2008, 2010) and globular (e.g. Hansen et al. 2004, 2007, Bedin et al.~2009) clusters. Methods to determine stellar population ages from their WD cooling sequences are usually based on the comparison of either the observed WD luminosity function (LF - star counts as a function of magnitude, e.g.,Winget et al.~1987, Bedin et al.~2010) or the actual bidimensional WD distribution in the CMD, with their theoretical counterparts (see, e.g., Hansen et al.~2007). Both techniques rely on an extensive use of grids of WD cooling sequences.  The sets of WD models largely employed in the more recent investigations on the age of WDs in Galactic stellar populations are those by Hansen~(1999, hereafter H99) and Salaris et al.~(2000, hereafter S00), computed with completely independent evolutionary codes and largely independent input physics. Additional recent large sets of WD evolutionary cooling models can be found in Althaus \\& Benvenuto~(1998 -- and later updates from the same group, who has also produced extensive libraries of He-core WD models, presented in Serenelli et al.~2002 and Althaus et al.~2009) and Fontaine, Brassard \\& Bergeron~(2001). The new grid of cooling models we present here to extend the evolutionary phase coverage of BaSTI, is an update of the results by S00. We include a complete set of both H- and He-atmosphere WD models (and isochrones, for a range of progenitor initial chemical compositions) that take advantage of the updated CO stratifications obtained from BaSTI AGB models, and employ boundary conditions from new sets of calculations of WD H and He atmospheres. Along with the presentation of our new models, we will discuss critically how WD cooling times are affected by the progenitor metallicity, uncertainties on the current estimate of the $^{12}C(\\alpha,\\gamma)^{16}O$ reaction rate and treatment of convection during the progenitor evolution. A similar analysis (albeit with several differences in the details) can be found in Prada-Moroni \\& Straniero~(2002), without taking into account CO phase separation upon crystallization in the WD cooling models.  The paper is structured as follows. Section~2 presents briefly the updates in the input physics compared to S00, and discusses critically our choices for the core chemical stratifications. Section~3 analyzes the main properties of the cooling models and WD isochrones, shows comparisons with S00 calculations and an example of application to study WDs in the solar neighborhood. A summary follows in Sect.~4.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % ",
        "conclusions": "\\label{summary}  We have expanded our BaSTI stellar evolution archive by including new, updated WD cooling models, computed using the CO stratification obtained from BaSTI AGB progenitor calculations. Improvements with respect to the S00 set of WD models concern the CO chemical profiles, that have been obtained employing an updated estimate of the $^{12}C(\\alpha,\\gamma)^{16}O$ reaction rate, and the inclusion of a full set of He-atmosphere WD models, computed with appropriate boundary conditions from non-gray model atmospheres. The reference set of WD models that will be made public at the BaSTI website makes use of the CO stratification at the first thermal pulse from progenitor models calculated with intial metal mass fraction Z=0.0198, and the inclusion of convective core overshooting during the main sequence. To assess how sensitive the models are to these assumptions, we have tested the effect of uncertainties on the recent determination of the $^{12}C(\\alpha,\\gamma)^{16}O$ reaction rate employed in the progenitor models, the inclusion/exclusion of core convective overshooting during the main sequence, different approaches for quenching the breathing pulses at the end of core He-burning, a variation of the metallicity of the progenitor, a variation of the number of pulses experienced by the progenitor models.  The results of this analysis indicate that the uncertainty on the $^{12}C(\\alpha,\\gamma)^{16}O$ reaction rate and the numerical approach used for inhibiting the breathing pulses have on the whole the largest impact on the WD cooling times, of about 7\\% at most- or of about 3\\% when the effect of phase separation upon crystallization is neglected. The progenitor metallicity, convective core overshooting during the main sequence phase, number of pulses before the WD formation, have overall a smaller effect.  We have discussed quantitatively differences in the mass-radius relationships and cooling speed of H- and He-atmosphere cooling models. The radii of the He-atmosphere models of a given mass are systematically lower than their H-atmosphere counterparts. Differences range between $\\sim$ 9\\% and $\\sim$ 3\\%, increasing with decreasing $M_{WD}$ and/or increasing temperature. He-atmosphere models show typically longer cooling times down to log($L/L_{\\odot}$)$\\approx -$4, before starting to cool down much faster at lower luminosities. We have also estimated the differences between ages of star clusters obtained employing our new H- and He-atmosphere WD models, as well as the S00 H-atmosphere WD calculations. Ages derived from S00 H-atmosphere models show only relatively small differences when compared to our new calculations.  As an example of application of our new set of models to real data, we have estimated an age of $\\sim$ 12~Gyr for the onset of star formation in the solar neighborhood, by fitting the local WD LF compiled by Catalan et al.~(2008). We have also studied the variation of the number ratio (N(He)/N(H)) with $T_{eff}$, predicted by our simulation of the local WDs. Due to the different cooling times of H- and He-atmosphere models, we show how this ratio changes with $T_{eff}$, increasing below $T_{eff}\\sim$10000~K, as observed. However, at least with our assumptions about the formation of the local WDs -- a constant progenitor formation rate and a (N(He)/N(H)) ratio at the onset of the WD phase that is constant with time -- the predicted ratio drops well below the observed value when $T_{eff}$ is lower than 7000-8000~K. This result can be explained in terms of the spectral transformation of a fraction of H-atmosphere objects, that increases with decreasing $T_{eff}$ below 7000-8000~K, reaching a maximum of $\\sim$24\\% at the lowest temperatures sampled by the observational data. As a consequence, one needs a broad range of H-layer thickness in solar neighborhood H-atmosphere WDs to explain these spectral changes, thicker envelopes being mixed with the underlying more massive He-layers at increasingly lower $T_{eff}$.  All cooling tracks and the reference chemical stratifications will be made publicly available at the official BaSTI website (\\url{http://www.oa-teramo.inaf.it/BASTI}). In addition, we provide WD isochrones for ages between 200~Myr and 14~Gyr for both H- and He-atmosphere objects (with and without the inclusion of phase separation) using as a reference the IFMR by Salaris et al.~(2009) and the progenitor lifetimes from BaSTI models including convective core overshooting on the main sequence. The isochrones will be available for progenitors with both scaled solar and $\\alpha$-enhanced mixtures, and 11 values of the metal fraction Z, ranging from Z=0.0001 to Z=0.04. For both cooling tracks and isochrones we provide magnitudes in the UBVRIJHK, and HST ACS photometric systems.  % %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1005/1005.5398_arXiv.txt": {
        "abstract": "Cygnus X-1 was the first X-ray source widely accepted to be a black hole candidate and remains among the most studied astronomical objects in its class. The detection of non-thermal radio, hard X-rays and gamma rays reveals the fact that this kind of objects are capable of accelerating particles up to very high energies.  In order to explain the electromagnetic emission from Cygnus X-1 in the low-hard state we present a model of a black hole corona with both relativistic lepton and hadron content. We characterize the corona as a two-temperature hot plasma plus a mixed non-thermal population in which energetic particles interact with magnetic, photon and matter fields. Our calculations include the radiation emitted by secondary particles (pions, muons and electron/positron pairs). Finally, we take into account the effects of photon absorption. We compare the results obtained from our model with data of Cygnus X-1 obtained by the COMPTEL instrument. ",
        "introduction": "The low-hard state of accreting black holes is characterized by the presence of a hot corona around the compact object. Figure \\ref{fig:geometria} shows a scheme of the main components of the system. For this geometry, we assume a spherical corona with a radius $R_{\\rm{c}}$ and an accretion disk that penetrates the corona up to a radius $R_{\\rm{d}}<R_{\\rm{c}}$. We suppose that the corona is homogeneous and in steady state.  \\begin{figure}[!h] \\centering \\includegraphics[clip,width=0.5\\textwidth, keepaspectratio]{Vieyro_1.eps} \\caption{Schematic representation of the components of the system. (Not to scale)} \\label{fig:geometria} \\end{figure}  We assume that the luminosity of the corona is 1 \\% of the Eddington luminosity of the black hole. For a corona characterized by the parameters of Table \\ref{tab_1}, this yields $L_{\\rm{c}}=1.3 \\times 10^{37}$ erg s$^{-1}$. The corona is composed of a two-temperature plasma, with an electron temperature $T_{e}=10^{9}$ K and an ion temperature $T_{i}=10^{12}$ K (e.g. Ref. \\refcite{Narayan}). This is a consequence of electron cooling being more efficient than proton.  In the low-hard state we can assume the system is in a steady state. We then consider equipartition of energy between the different components of the system and obtain the values of the relevant physical parameters.  The X-rays emission of the corona is characterized by a power law with an exponential cut-off at high energies, \\begin{displaymath} n_{\\rm{ph}}(E) = \\left\\{ \\begin{array}{lrl} A_{\\rm{ph}} E^{-\\alpha} &  E & < E_{\\rm{c}}\\\\ A_{\\rm{ph}} E^{-\\alpha}  e^{-E/E_{\\rm{c}}}  &   E & \\geq E_{\\rm{c}}. \\end{array} \\right. \\end{displaymath}  \\noindent In accordance with what is observed in Cygnus X-1 (e.g. Ref. \\refcite{Poutanen}), we adopt $\\alpha=1.6$ and $E_{\\rm{c}}=150$ keV.  Table \\ref{tab_1} shows the values of the parameters in our model.  \\begin{table}[!h] \\tbl{Model parameters.} {\\begin{tabular}{@{}lc@{}} \\toprule  Parameter & Value\\\\[0.01cm] \\hline\\noalign{\\smallskip} $M_{\\rm{BH}}$:  black hole mass [$M_{\\odot}$]\t\t\t\t\t\t\t& $10$$^{(1)}$\t\t\t\t\t\\\\[0.01cm] $R_{\\rm{c}}$:   corona radius [cm] \t\t\t\t\t\t\t\t\t\t \t  & $5.17 \\times 10^{7}$$^{(1,2)}$ \t\\\\[0.01cm] $T_{e}$:        electron temperature [K] \t\t\t\t\t\t\t\t\t&\t$10^9$\t\t\t\t\t\t\t\t\\\\[0.01cm] $T_{i}$:        ion temperature [K] \t\t\t\t\t\t\t\t\t\t\t&\t$10^{12}$ \t\t\t\t\t\t\\\\[0.01cm] $E_{\\rm{c}}$:   X-ray spectrum cut-off [keV]\t\t\t\t\t\t\t& $150$       \t\t\t\t\t\\\\[0.01cm] $\\alpha$: \t\t\tX-ray spectrum power-law index    \t\t\t\t& $1.6$\t\t\t\t\t\t\t\t\t\\\\[0.01cm] $\\eta$: \t\t\t\tacceleration efficiency \t\t\t\t\t\t\t\t\t& $10^{-2}$\t\t\t\t\t\t\t\\\\[0.01cm] $B_{\\rm{c}}$: \tmagnetic field [G]\t\t\t\t \t\t\t\t\t\t\t\t& $2 \\times 10^7$\t\t\t\t\\\\[0.01cm] $n_{i},n{e}$:   plasma density [cm$^{-3}$] \t\t\t\t\t\t\t\t& $6.2 \\times 10^{13}$\t\\\\[0.01cm] $a$: \t\t\t\t\t\thadron-to-lepton energy ratio \t\t\t\t\t\t& $1$ - $100$\t\t\t\t\t\t\\\\[0.01cm] $kT$:\t\t\t\t\t\tdisk characteristic temperature [keV] \t  & $0.1$\t\t\t\t\t\t\t\t\t\\\\[0.01cm] $v$:\t\t\t\t\t\tadvection velocity  [c] \t \t\t& $0.1$\t\t\t\t\t\t\t\t\t\\\\[0.01cm]  \\hline\\\\[0.005cm] \\multicolumn{2}{l}{ $^{(1)}$ Typical value for Cygnus X-1 (Ref. \\refcite{Poutanen}).}\t\t\t\t  \\\\[0.01cm] \\multicolumn{2}{l}{ $^{(2)}$ $35 R_{\\rm{G}}$.} \t\t\t \t\t\t\t\t\t\t\t\t\t\t\t\t\t\t\t\t\t\t\t\\\\[0.01cm] \\end{tabular} \\label{tab_1}} \\end{table} ",
        "conclusions": "On one hand, the SED predicted by the model with an injection dominated by protons and diffusion reproduces the observations of Cygnus X-1. On the other hand, our model predicts a high energy emission, above 1 TeV, that might be detectable by instruments like MAGIC II or CTA in the near future."
    },
    "1005/1005.0097_arXiv.txt": {
        "abstract": "Despite our present-day inability to predict the topology of the universe it is expected that we should be able to detect it in the near future. A nontrivial detectable topology of the space section of the universe can be probed for all homogeneous and isotropic universes through the circles-in-the-sky. We discuss briefly how one can use this observable attribute to set constraints on the dark energy equation of state parameters. ",
        "introduction": "In standard  cosmology the Universe is described by a $4$--manifold $\\mathcal{M} = \\mathbb{R}\\times M_3$ endowed with the spatially homogeneous and isotropic Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) metric. Furthermore, the spatial sections $M$ are often assumed to be the simply-connected $3$--manifolds: Euclidean $\\mathbb{E}^3$, spherical $\\mathbb{S}^3$, or hyperbolic space $\\mathbb{H}^3$. However, the majority of constant curvature $3$--spaces, $M_3$,  are multiply-connected quotient manifolds of the form $\\mathbb{R}^3/\\Gamma$, $\\mathbb{S}^3/\\Gamma$, and $\\mathbb{H}^3/\\Gamma$, where $\\Gamma$ is fixed-point free group of isometries (see, e.g., the review Refs.~\\refcite{CosmTopReviews} for details). An important observational consequence of a nontrivial (multiply-connected) observable spatial non-trivial topology\\cite{TopDetec} of $M_3$ is the existence of the circles-in-the-sky,\\cite{CSS1998} i.e. pairs of matching circles with the same distribution of temperature fluctuations, identified by $\\Gamma$. Hence, to observationally probe a putative non-trivial topology of $M_3$, one ought to extract the pairs of correlated circles from full-sky CMBR maps.  Here we complement concisely our previous works\\cite{Previous} by showing how a possible detection of a pair of circles-in-the-sky can be used to set constraints on the dark energy equation of state (EOS) parameters in \\emph{globally homogeneous} universes.  \\vspace{-0.3cm} ",
        "conclusions": ""
    },
    "1005/1005.5167_arXiv.txt": {
        "abstract": "{Massive evolved stars can produce large amounts of dust, and far-infrared (IR) data are essential for determining the contribution of cold dust to the total dust mass. Using $Herschel$, we search for cold dust in three very dusty massive evolved stars in the Large Magellanic Cloud: R71 is a Luminous Blue Variable, HD\\,36402 is a Wolf-Rayet triple system, and IRAS05280--6910 is a red supergiant. We model the spectral energy distributions using radiative transfer codes and find that these three stars have mass-loss rates up to $10^{-3}~M_\\odot~{\\rm yr}^{-1}$, suggesting that high-mass stars are important contributors to the life-cycle of dust. We found far-IR excesses in two objects, but these excesses appear to be associated with ISM and star-forming regions. Cold dust ($T<100$~K) may thus not be an important contributor to the dust masses of evolved stars.} ",
        "introduction": "Intermediate-mass asymptotic giant branch (AGB) stars are potentially the dominant dust source in the Galaxy \\citep{gehrz89} and in low-metallicity environments like the Large Magellanic Cloud \\citep[LMC;][]{srinivasan09, matsuura09} and other dwarf galaxies \\citep{boyer09_dirr}. However, dust production in high-mass stars ($\\gtrsim$$8~M_\\odot$) remains uncertain. It has been suggested that supernovae (SNe) might be the dominant dust factory at high-$z$, since intermediate-mass stars have not yet had time to evolve into AGB stars \\citep{morgan03, dwek09}. However, the amount of dust formed in SNe in the local universe is much less than required to explain the dust seen at high-$z$ \\citep[e.g.,][]{sugerman06,andrews10}. It is also unclear if the dust forms before or after the SN explosion.  Alternatively, \\citet{sloan09} and \\citet{valiante09} show that AGB stars {\\it can} contribute dust at high redshifts.  These studies point to a need to measure the total dust mass from all types of stars to obtain a global picture of dust evolution in galaxies.  Part of the LMC was observed with {\\it Herschel} \\citep[][this volume]{pilbratt} as part of the science demonstration program (SDP) and the Legacy program entitled {\\it A HERschel Inventory of The Agents of Galaxy Evolution} \\citep[HERITAGE;][this volume]{meixner}. In this letter, we describe a first look at {\\it Herschel} Photodetector Array Camera and Spectrometer \\citep[PACS;][this volume]{poglitsch} and Spectral and Photometric Imaging REceiver \\citep[SPIRE;][this volume]{griffin} detections of 3 examples of dust-producing massive evolved stars in the LMC: the Wolf-Rayet (WR) system HD\\,36402, the luminous blue variable (LBV) HDE\\,269006 (or R71), and the red supergiant (RSG) OH/IR star IRAS05280--6910. While such stars have been studied extensively in the mid-IR \\citep[e.g.,][]{morrisiso99, clark03, crowther07, bonanos09, vanloon10}, this study is among the first to probe them at $\\lambda\\gtrsim200~\\micron{}$.   \\begin{table*}% \\caption{Target Information and {\\it Herschel} Flux Densities} \\label{tab:targets} \\centering \\begin{tabular}{llllllll}% \\hline \\hline Name & R.A., Dec. (J2000) & Type & \\multicolumn{5}{c}{$F_\\nu$ (mJy)}\\\\ & & & 100~\\micron & 160~\\micron & 250~\\micron & 350~\\micron & 500~\\micron \\\\ \\hline R71 & 05$^{\\rm h}$02$^{\\rm m}$07.39$^{\\rm s}$, $-$71${\\degr}$20${\\arcmin}$13.1${\\arcsec}$  & LBV & $447\\pm8$ & $105\\pm10$ & $50\\pm6$ & $49\\pm13$ & $27\\pm9$\\\\ IRAS05280$-$6910 & 05$^{\\rm h}$27$^{\\rm m}$40.11$^{\\rm s}$, $-$69${\\degr}$08${\\arcmin}$04.5${\\arcsec}$ & RSG, OH/IR &...&...& $205\\pm8$ & $46\\pm13$ & $<$$11\\pm10$ \\\\ HD\\,36402 & 05$^{\\rm h}$26$^{\\rm m}$03.94$^{\\rm s}$, $-$67${\\degr}$29${\\arcmin}$57.9${\\arcsec}$ & WC4($+$O?)$+$O8I &...&...&...&...&...\\\\ HD\\,36402 IR1/YSO2& 05$^{\\rm h}$26$^{\\rm m}$02.99$^{\\rm s}$, $-$67${\\degr}$29${\\arcmin}$57.7${\\arcsec}$ &YSO?& $531\\pm10$ & $872\\pm33$ & $677\\pm33$ & $<$$431\\pm23$ & $<$$132\\pm21$\\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[{\\sc Notes.}--] Quoted uncertainties are 1\\,$\\sigma$. Upper limits are either below the 3\\,$\\sigma$ detection level or include substantial flux from an adjacent unresolved source. IRAS05280--6910 does not fall within the PACS coverage.  The WR star itself is not resolved in {\\it Herschel} images; it is severely blended with an IR source (IR1/YSO2).  The 350 and 500~\\micron{} fluxes of IR1/YSO2 include a nearby YSO (N51-YSO1), which is unresolved from IR1/YSO2 at these wavelengths. Aperture corrections for PACS and SPIRE were estimated using the current point-spread functions, and are roughly 1.4 and 1.2, respectively. None of the sources are heavily affect by high and/or variable backgrounds (Fig.~\\ref{fig:raw}). \\end{list} \\end{table*}  \\begin{figure} \\centering \\hbox{ \\includegraphics[width=0.155\\textwidth]{14513fg1a.eps} \\includegraphics[width=0.155\\textwidth]{14513fg1b.eps} \\includegraphics[width=0.155\\textwidth]{14513fg1c.eps} } \\caption{{\\it Herschel} images of R71 (left panels), IRAS05280-6910 (middle panels), and HD\\,36402 (right panels, also see Fig.~\\ref{fig:img}). Contours on a linear scale are included where it is difficult to see the detection. \\label{fig:raw}} \\end{figure} ",
        "conclusions": ""
    },
    "1005/1005.3339_arXiv.txt": {
        "abstract": "We describe the cryostat and supporting electronics for the EBEX experiment. EBEX is a balloon-borne polarimeter designed to measure the B-mode polarization of the cosmic microwave background radiation. The instrument includes a $1.5$ meter Gregorian-type telescope and 1432 bolometric transition edge sensor detectors operating at 0.3 K. Electronics for monitoring temperatures and controlling cryostat refrigerators is read out over CANbus. A timing system ensures the data from all subsystems is accurately synchronized. EBEX completed an engineering test flight in June 2009 during which the cryogenics and supporting electronics performed according to predictions. The temperatures of the cryostat were stable, and an analysis of a subset of the data finds no scan synchronous signal in the cryostat temperatures. Preparations are underway for an Antarctic flight. ",
        "introduction": "\\label{sec:Science} EBEX (E and B EXperiment) is a balloon-borne experiment equipped with 1432 bolometric transition edge sensors (TES). It is designed to measure the temperature and polarization of the cosmic microwave background (CMB) radiation. EBEX has several scientific goals. The first is to search for the B-mode polarization of the CMB which is a signature of inflationary gravitational waves [\\citen{kamionkowski97b},\\citen{seljak97}]. The second goal is to characterize Galactic polarized dust emission and to determine its angular power spectra in both E and B-mode polarizations at the micro-Kelvin level. Characterization of dust foregrounds is a key element in extracting the B-mode signal [\\citen{brandt94},\\citen{baccigalupi03}]. The third goal is to measure the predicted but yet undetected lensing of the CMB. Gravitational lensing of the CMB distorts the anisotropy pattern on the sky, modifies its power spectrum [\\citen{seljak97a}], and converts E-mode to B-mode polarization [\\citen{zaldarriaga98}].  Previous EBEX related publications describe the experiment's science goals [\\citen{Oxley04}], design [\\citen{granger08}] and detectors and readout [\\citen{hubmayr08}]. Here we will focus on the cryostat and its housekeeping readout. \\begin{figure} \\label{fig:EBEX_payload} \\begin{center} \\parbox{2.1in}{\\epsfig{file=ebex_hanging.eps,width=2in }} \\hspace*{4pt} \\parbox{2.1in}{\\epsfig{file=cryostat_cross_rr.eps,width=3in}} \\caption{Left: The EBEX payload before launch. Right: A cross section of the EBEX cryostat.} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "1005/1005.4802_arXiv.txt": {
        "abstract": "{The understanding of the origin of dark matter has great importance for cosmology  and particle physics. Several interesting extensions of the standard model dealing with solution of this problem motivate the concept of hidden sectors consisting of $SU(3)_C\\times SU(2)_L\\times U(1)_Y$ singlet fields. Among these models, the mirror  matter model is certainly one of the most interesting. The model explains the origin of parity violation in weak interactions, it could also explain the baryon asymmetry of the Universe and provide a natural ground for the explanation of dark matter. The mirror matter could have  a portal to our world through photon-mirror photon  mixing ($\\epsilon$). This mixing would lead to orthopositronium  ($o-Ps$) to mirror orthopositronium oscillations, the experimental signature of which is  the apparently invisible decay of $o-Ps$. In this paper, we describe  an experiment  to search for the decay  $o-Ps\\to invisible $ in vacuum  by using a pulsed slow positron beam and a massive 4$\\pi$ BGO crystal calorimeter. The developed high efficiency positron tagging system, the low calorimeter energy threshold and high hermiticity allow the expected sensitivity in mixing strength to be $\\epsilon \\simeq 10^{-9}$, which is more than one order of magnitude below the current Big Bang Nucleosynthesis limit and in a region of parameter space of great theoretical and phenomenological interest. The vacuum experiment with such sensitivity is particularly timely in light of the recent DAMA/LIBRA observations of the annual modulation signal consistent with a mirror type dark matter interpretation. } ",
        "introduction": "Cosmological observations of  galactic rotational curves \\cite{Borriello:2000rv} and the gravitational lensing \\cite{Hoekstra:2002nf,Metcalf:2003sz} give strong evidence for the existence of dark matter  (see e.g. \\cite{Bertone:2004pz} for an excellent review). In addition, the recent confirmation by the DAMA/LIBRA  experiment \\cite{DAMALIBRA} of the annual modulation  signal observed  by the DAMA/NaI \\cite{dama} seems  provide the first direct terrestrial experimental  observation of the existence of non-baryonic dark matter in our galactic halo. The need to explain these hints of the presence of dark matter provides one of the strongest indications for the existence of new physics beyond the Standard Model (SM) and the identification of the origin of dark matter is a task of enormous importance for both particle physics and cosmology. At present, the most popular candidates for the (thermal-produced) dark matter are the so-called weakly interacting massive particles (WIMPs), which are e.g. lightest supersymmetric particles, Kaluza-Klein particles in universal extra dimension models, axions etc... However, despite of significant efforts the experiments looking for WIMPs lead so far to negative results, thus, pushing further possible WIMP searches  into  a very high-energy and/or high sensitivity frontiers, for a recent review see e.g. \\cite{feng} and references therein.  An additional natural ground for the explanation of dark matter is provided by a class of interesting  theoretical models introducing the concept of ''hidden'' sectors consisting of $SU(3)_C\\times SU(2)_L\\times U(1)_Y$ singlet fields. It is worthwhile to note that, even in the SM, some of the matter fields are singlet under one or more of the colour and electroweak gauge groups. Thus, the extension to include a further sector which transforms under the new but not under the familiar gauge symmetries is not particularly exotic from a theoretical viewpoint.  The sensitivity of searches for the new singlet particles depends in detail on their  couplings and  mass scale, e.g. if the mass scale of a hidden sector is too high, it is experimentally unobservable and indeed is hidden. Then the question arises: could the important sensitive  searches for the hidden sectors be performed at low energy frontier ? The answer for this question is definitely positive. For example, there is a class of models with at least one additional $U(1)$ gauge factor where the corresponding hidden gauge boson could be light, or even massless \\cite{mfok, holdom, mf}, for a recent review see  \\cite{jr}. With respect to the mass and lifetime range for hidden sector states, the niche with masses in the range of a few MeV to a few GeV, with lifetimes of less than 1 second, is also quite intriguing as many of the most severe astrophysical and cosmological constraints that apply to lighter states are weakened or eliminated, while those from high-energy colliders are often inapplicable \\cite{pospel}.  Recent hints from terrestrial and  astrophysical anomalies, in particular first CDMSII result \\cite{cdms} and annual modulations from DAMA/LIBRA at Gran Sasso \\cite{DAMALIBRA, dama}, positron and/or electron-positron pair excess in cosmic ray from PAMELA \\cite{pam} and ATIC \\cite{atic} (however, the ATIC excess seems has not been confirmed by Fermi-LAT \\cite{fermi}) are found to be consistent with an interpretation in context of dark matter charged under new gauge fields \\cite{dmreview}. Thus, if those observations are really signals for DM they would indicate particularly exotic forms of DM. This suggests  that testing hidden sectors with the relatively light mass scale states in the MeV-to-GeV  range  weakly coupled to the SM model represents an intriguing possibility.  Then it is important to ask  which hidden sector models  could be tested with the best experimental sensitivity? Certain classes of hidden sectors may be  tested, but it is clear that not  many generic scenarios, which are viable, could be strongly  constrained and even excluded on the base of precision measurements. Among numerous  candidates that have been discussed, one of the most promising, which could reconcile the  DAMA annual modulation signal with the results of the CDMS II experiment, is mirror type dark matter \\cite{foot4,foot5}. In addition, the CoGeNT collaboration has recently reported a rising low energy spectrum in their ultra low noise germanium detector \\cite{cogen}. This is particularly interesting as the energy range probed by CoGeNT overlaps with the energy region in which DAMA has observed their annual modulation signal. It has been recently shown \\cite{foot6}, that the mirror dark matter candidate can simultaneously explain both the DAMA annual modulation signal and the rising low energy spectrum observed by CoGeNT. This constitutes a model dependent confirmation of the DAMA signal and adds weight to the mirror dark matter paradigm.  In this paper, we show that the proposed search could result either in a strong evidence of the existence of dark matter of the mirror type, or completely exclude this type of new physics beyond the standard model. It should be noted that mirror baryons are naturally dark, stable and massive. Currently, it seems that this type of matter could also explain in a natural way the coincidence between visible and dark matter densities in the universe ($\\Omega_B=0.044$ and $\\Omega_{DM}=0.26$) \\cite{berezhiani,Ciarcelluti:2008vs}. Interestingly, it could also provide an explanation of the controversial behaviour in galaxy cluster collisions \\cite{silagadze}. Furthermore, if mirror matter is present in our universe it would mean that Parity (spatial-inversion) is an unbroken symmetry of nature. This gives an exciting motivation for testing this model by a laboratory experiment at low energies.  The rest of the paper is organized as follows. In Sec. \\ref{sec:mirror}, we review the mirror model and work out briefly the mirror Higgs boson physics at the CERN LHC, and in more details, the effect of oscillation of ordinary positronium to the mirror one and its experimental consequences.  We also present the calculations we performed in order to estimate the effect of matter and external fields on the oscillation probability. In Section \\ref{sec:exptech}, we present the experimental technique. In Section \\ref{sec:setup}, the design of the experiment and the detector components are described in detail. The simulations of the signal and background sources, as well as the expected sensitivity  are discussed in Section \\ref{sec:Bkg_vac} and \\ref{sec:sensitivity}, respectively. Section \\ref{sec:conclusion} contains concluding remarks. ",
        "conclusions": "\\label{sec:conclusion}  In this paper, a proposed search for mirror-type dark matter, looking for \\ops\\ invisible decays in a vacuum cavity, is presented. In the Introduction,  the mirror matter relevance to the dark matter problem and its link with positronium has been reviewed. The effect of external fields (electric and magnetic) on the oscillation probability and of the collisions of \\ops\\ with the cavity walls were estimated. In Section \\ref{sec:setup}, the design and the experimental results for the different parts of the experiment have been presented. Section \\ref{sec:Bkg_vac} includes the simulation results of the background estimation and Section \\ref{sec:sensitivity} the estimated sensitivity. The goal is to reach a sensitivity in the branching ratio of $Br(\\ops\\to invisible)\\simeq 10^{-7}$ to confront the annual modulation signal observed by DAMA/NaI and DAMA/LIBRA (with 8.2$\\sigma$ significance)  with mirror dark matter scenarios. In case of a signal observation, the experiment would offer a unique and essential feature, whereby an increase or decrease of the signal rate by a factor $\\sim 2$  is possible while keeping the background constant. In case of a null result, this search will provide a limit on the photon-mirror photon mixing strength about one order of magnitude better than presently derived from Big Bang Nucleosynthesis.     \\vskip0.5cm   {\\bf Acknowledgments}  We wish to thank Z. Berezhiani, R. Eichler, R. Foot, D. S. Gorbunov, S. Karshenboim, N. V. Krasnikov, C. Lenz-Cesar, V. Matveev,  V. A. Rubakov, and  A. A. Tokareva for the very useful discussions. We acknowledge A. Antognini and F. Kottmann for lending us some carbon foils. We are grateful to A. Badertscher, N. A. Golubev, L. Knecht, L. Liszkay, J. P. Peigneux, P. Perez, D. Sillou for their essential help. This work was supported by the ETH Zurich, by the Swiss National Science Foundation, by the Institute for Nuclear Research of the Russian Academy of Sciences and by the CNPq (Brazil). We thank the CERN for its hospitality."
    },
    "1005/1005.1420_arXiv.txt": {
        "abstract": "We present a family of dynamic rotating cylindrically symmetric  Ricci-flat  gravitational fields  whose  geodesic motions have the structure of  gravitomagnetic jets. These correspond to helical motions of free test particles up and down parallel to the axis of cylindrical symmetry and are reminiscent of the motion of test charges in a magnetic field. The speed of a test particle in a gravitomagnetic jet asymptotically approaches the speed of light. Moreover, numerical evidence suggests that jets are attractors.  The possible implications of our results for the role of gravitomagnetism in the formation of astrophysical jets are briefly discussed. ",
        "introduction": "A constant uniform \\emph{magnetic} field configuration in an inertial frame of reference has cylindrical symmetry; therefore, the motion of a test charge in this field is such that the particle's momentum and angular momentum in the direction of the field are constants of the motion. The particle in general moves with constant speed on a helix whose axis is along the field direction; moreover, the radius and step of the helical path are constants as well. The sense of helical motion about the direction of the magnetic field is positive (negative) for a test particle of negative (positive) electric charge. If the initial velocity of the particle is normal to the direction of the magnetic field, then the particle simply moves along a circle in the plane perpendicular to the field.  The purpose of this paper is to study the analogous situation for the motion of free test particles in a \\emph{gravitomagnetic} field. The nonlinearity of this field implies that only a rough similarity may be anticipated. We investigate geodesic  motion in a rotating dynamic spacetime region that is a Ricci-flat solution of Einstein's equations with cylindrical symmetry~\\cite{1}. Though this time-dependent gravitational case is considerably more complicated than the magnetic case, we find qualitatively similar phenomena. In fact, the helical motions up and down parallel to the axis of symmetry are reminiscent of the double-jet structure of certain high-energy astrophysical sources.  Quasars and active galactic nuclei generally exhibit distinct relativistic outflows. These jets are conjectured to originate from massive rotating black holes surrounded by accretion disks; the outflows are focused beams of relativistic particles that proceed up and down along the rotation axis of the black hole (see, for instance,~\\cite{n2}). Similar phenomena have been observed in other high-energy astrophysical sources such as the Galactic X-ray binary systems. While our results suggest a mechanism for astrophysical jet formation, the more complicated physical process involves general relativistic MHD~\\cite{2,f}. The present work---together with previous efforts~\\cite{3,4,5,6,7,9}---contributes to the purely gravitational aspects of this fundamental problem in astrophysics.  For a subclass of the Ricci-flat solutions under consideration, we show that the geodesic equations have families of special exact solutions that we call gravitomagnetic jets.  More precisely, a gravitomagnetic jet is a set of special geodesics. These  generally exhibit helical motions about the axis of  cylindrical symmetry and their union is a non-compact connected invariant manifold that attracts all nearby geodesics. While we highlight features of these jets in the source-free cylindrical spacetime region of interest and provide strong (numerical) evidence that these families are  attractors, the interesting question of the nature of the external matter currents that could generate such a gravitational field remains beyond the scope of our present investigation.  According to general relativity, a rotating mass generates a relativistic, and hence non-Newtonian, gravitomagnetic field that is due to mass current. The exterior gravitomagnetic field of the Earth has recently been \\emph{directly} measured via Gravity Probe B (GP-B)~\\cite{13}. On the theoretical side, solutions of Einstein's equations in the case of axial symmetry have received attention for a long time (see Ch. VIII of~\\cite{16} and references cited therein). In particular, rotating solutions with cylindrical symmetry have been investigated by a number of authors (see~\\cite{14,11n,15} and references therein).  Previous interesting work on exact cylindrically symmetric gravitational fields in connection with the origin and structure of astrophysical jets has mainly involved the study of geodesics in the interior of time-independent rigidly rotating dust cylinders~\\cite{13n}. The behavior of free test particles in this case is directly influenced by the gravitational attraction of the rotating dust particles; to avoid this circumstance, we concentrate here on certain source-free gravitational fields that happen to depend exponentially upon time. These Ricci-flat fields are perhaps more representative of the strongly time-dependent near-zone exteriors of accreting and growing gravitationally collapsed configurations where astrophysical jets are expected to originate. The temporal variation of the gravitational fields considered in our work leads to a significant and surprising feature of the special exact solutions of the geodesic equations. The speeds of test particles in gravitomagnetic jets start from values that are always above a certain minimum speed, rapidly increase along their paths and asymptotically approach the speed of light. The minimum speed represents the speed of circular motion perpendicular to the axis of rotation.  The range of this minimum speed turns out to be from zero up to about $0.63\\, c$.  Is the main general relativistic problem of high-energy astrophysical jets solved in our paper? Our gravitational model is certainly too elementary to be an adequate representation of the complex physical situation. Nevertheless, in our simple model free test particles appear to be exponentially accelerated to almost the speed of light resulting in ultrarelativistic jet streams parallel and antiparallel to the axis of rotation.  The plan of the paper is as follows. The class of spacetime metrics that we study in this paper is described in section II. Each metric in this class involves a function $X (r)$ that is a solution of an ordinary differential equation and must be so chosen as to render our system of cylindrical coordinates admissible in the spacetime region of interest. The procedure for the determination of an appropriate $X (r)$  is described in section III. The rotational aspects of the resulting gravitational field are discussed in section IV.  Section V treats the motion of free test particles and null rays in the spacetime region of interest. This section also contains a discussion of the special analytic solutions of the geodesic equations that form  gravitomagnetic jets. Numerical evidence that gravitomagnetic jets are indeed attractors is presented in section VI. Section VII contains a discussion of our results. For clarity of presentation, some of the detailed calculations and mathematical arguments are relegated to the appendices.  Specifically, Appendix~\\ref{appen:A} contains useful formulas related to the principal dynamic spacetime metric under consideration in this paper. The proofs of the main mathematical results regarding the solutions of a certain highly nonlinear ordinary differential equation are given in Appendix~\\ref{appen:B}, where, due to the nature of the subject matter, some of the notation employed is independent of the rest of the paper. Appendix~\\ref{appen:C} treats the geodesics of the time-reversed metric. ",
        "conclusions": "The cylindrical symmetry employed throughout this work has been a useful simplifying assumption. However, axial symmetry is expected to be a better approximation for the treatment of high-energy astrophysical jets that appear in circumstances involving a rotating collapsed configuration surrounded by a rotating accretion disk. On the other hand, in this case an analytical treatment appears to be prohibitively complicated.  How could one generate the gravitational fields discussed in this paper? In electrodynamics, the magnetic field $\\mathcal{B}$ inside an infinite circular cylinder is uniform and parallel to the axis of symmetry, provided the cylinder is surrounded by a uniform current sheet. The connection between the source and the interior field is given by $ \\mathcal{B} = 4 \\pi i /c$, where $i$ is the amount of electric current per unit length of the cylinder. This is a good approximation around the axis near the center of a long solenoid. In our gravitational case, we would expect that the source-free dynamic interior solution could be joined---perhaps along the inner and outer boundaries---to an exterior solution that could serve as the source of the interior field.  However, finding such a source requires a separate investigation that is beyond the scope of this paper. In fact, it may be advantageous to look for gravitomagnetic jets in more realistic axisymmetric systems in general relativity.  Despite these drawbacks, it is remarkable that general relativity permits the existence of Ricci-flat rotating cylindrical gravitational fields that admit gravitomagnetic jets, which are formed from solutions of the geodesic equation corresponding to generally helical motions of free test particles up and down parallel to the axis of symmetry with speeds that asymptotically approach the speed of light.  Particle acceleration mechanisms are important in astrophysics~\\cite{vn24}. We have shown that general relativity can in principle provide a purely gravitational mechanism for the directional acceleration of test particles to ultrarelativistic speeds.           \\appendix \\setcounter{section}{0} \\newcounter{saveeqn}% \\newcommand{\\alpheqn}{\\setcounter{saveeqn}{\\value{equation}} \\stepcounter{saveeqn}\\setcounter{equation}{0} \\renewcommand{\\theequation}{\\mbox{\\Alph{section}\\arabic{equation}}}} \\renewcommand{\\thesection}{\\Alph{section}}"
    },
    "1005/1005.3755_arXiv.txt": {
        "abstract": "{We analysed the \\xmm\\ archival observations of 16 neutron star (NS) low-mass X-ray binaries (LMXBs) to study the Fe~K emission in these objects. The sample includes all the observations of NS LMXBs performed in EPIC pn Timing mode with \\xmm\\ publicly available until September 30, 2009. We performed a detailed data analysis considering pile-up and background effects. The properties of the iron lines differed from previous published analyses due to either incorrect pile-up corrections or different continuum parameterization. 80\\% of the observations for which a spectrum can be extracted showed significant Fe line emission. We found an average line centroid of 6.67\\,$\\pm$\\,0.02~keV and a finite width, $\\sigma$, of 0.33\\,$\\pm$\\,0.02~keV. The equivalent width of the lines varied between 17 and 189 eV, with an average weighted value of 42\\,$\\pm$\\,3~eV. For sources where several observations were available the Fe line parameters changed between observations whenever the continuum changed significantly. The line parameters did not show any correlation with luminosity. Most important, we could fit the Fe line with a simple Gaussian component for all the sources. The lines did not show the asymmetric profiles that were interpreted as an indication of relativistic effects in previous analyses of these LMXBs. ",
        "introduction": "\\label{sect:intro}  Accreting binaries often show iron line emission in their X-ray spectra. The ability of \\chandra, \\xmm, \\suzaku\\ and {\\it Swift} to obtain medium to high resolution spectra covering the iron K energy band has opened a new era in the observations of stellar-mass black holes (BHs) and NS binaries. The large effective area of these observatories is crucial for the detection of iron line emission while the high resolution enables us to discriminate between narrow and intrinsically broad features. Broad iron lines are also present in Active Galactic Nuclei \\citep{tanaka95nat, nandra97apj} though studies of large samples of sources show that the average fraction of broad Fe lines is never higher than 40\\% \\citep{guainazzi06an, nandra07mnras}.  The origin of broad iron lines has been extensively discussed in the literature \\citep[e.g.][]{sunyaev80aa, cygx1:fabian89mnras, done07mnras, ross07mnras, titarchuk09apj}. However, even after the advent of the current powerful X-ray observatories, the exact determination of the line width and the mechanisms responsible for it are still controversial. Undoubtedly the most exciting possibility is that such lines originate in the disc close to the BH event horizon or to the NS by fluorescent emission following illumination (and photoinisation) of the accretion disc by an external source of X-rays \\citep[][]{reynolds03ps, fabian05, matt06an}. In this model, the combination of relativistic Doppler effects arising from the high orbital velocities and gravitational effects due to the strong gravitational field in the vicinity of the compact object smear the reflected features. In this case, detailed X-ray spectroscopy of iron line features can be used to study Doppler and gravitational redshifts, thereby providing key information on the location and kinematics of the material in the vicinity of the compact object. Most interesting of all is the potential for establishing BH spin using relativistic iron \\ka\\ lines. The spin value is constrained mainly by the lower boundary of the broad line, which depends on the inner boundary of the disc emission (identified with the marginally stable orbit) where the gravitational redshift is maximal. Therefore, lines which are strongly skewed toward lower energies can indicate black hole spin \\citep[][]{miller09apj, 1753:hiemstra09mnras, 1655:reis09mnras}.  The recent claim of broad skewed iron lines from NS binaries has opened the exciting possibility of determining an upper limit to the radius of the NS, the most difficult parameter to obtain in order to constrain the equation of state of NSs \\citep{serx1:bhattacharyya07apjl}. The advantage of using iron K emission lines as a probe of the NS radius is that they only require short observations to clearly reveal the relativistic lines, and do not require any knowledge of the distance to the object.  An alternative location of the line emitting region could be the inner part of the so-called Accretion Disc Corona (ADC), formed by evaporation of the outer layers of the disc illuminated by the emission of the central object \\citep[e.g.,][]{white82apj}. In the corona, where the gas is highly ionised, the iron line emission is likely to come from recombination onto hydrogenic and helium-like iron. In the region below the transition region the iron will be at most a few times ionised and fluorescence will be the dominant process \\citep{kallman89apj}.  In the corona, Compton scattering will broaden and shift the centroid of the line. The line broadening and shift are due to Doppler and/or recoil effects, depending on the temperature of the plasma compared to the energy of the photons \\citep{pozdnyakov79aa, sunyaev80aa}.  The final line width is determined by the simultaneous effects of blending, Compton scattering and rotation. For a value of {\\it f}\\,$\\sim$\\,0.1 ($f$ being the fraction of the incident X-ray flux which is assumed to penetrate the base of the hot medium or corona), \\citet{kallman89apj} calculated a centroid energy of 6.6--6.7~keV for the iron line and a width of $\\sim$0.6~keV at small radii (R~$\\approxlt$~10$^8$~cm).  For {\\it f}\\,$\\sim$\\,1, the width can be as large as $\\sim$1~keV (FWHM). Alternatively, widths of $\\sim$1~keV can also be attained via emission from a photoionised gas with large Thomson depth ($\\tau \\approxgt$3). Such large Compton depths could arise if the corona is heated by some mechanism other than X-rays, or if the line of sight of the observer through the corona were very different from that of the compact X-ray source \\citep{kallman89apj}.  A third scenario for the formation of broad iron lines in LMXBs could be one where extensive red wings form by recoil of line photons in an optically thick medium expanding, or converging, at relativistic velocities \\citep{titarchuk03apj, laming04apj, laurent07apj}. It is assumed that the wide-open wind is launched at some disc radius where presumably the local radiation force exceeds the local disc gravity. The wind should be illuminated by the emission of X-rays formed in the innermost part of the source. The \\ka\\ line is generated in a narrow wind shell and the line profile is formed in the partly ionised wind when \\ka\\ line photons are scattered off the diverging flow (wind) electrons \\citep{titarchuk09apj}.  The red-skewed part of the spectrum is formed by photons undergoing multiple scatterings while the primary peak is formed by photons escaping directly to the observer.  The line properties of an ADC are in general expected to be different from those of an accretion disc \\citep{brandt94mnras}. In the former case the equivalent width (\\ew) is nearly independent of inclination except at very high angles \\citep{vrtilek93apj}, while the dependence on inclination of the case of an accretion disc is significant \\citep{brandt94mnras}. Therefore iron lines appear to be a powerful tool for distinguishing between an ADC, an accretion disc, or any other origin of the iron line in LMXBs. In any case, a proper modelling of the iron line is mandatory to determine uniquely the origin of the line and, in the case of relativistic broadened iron lines, to derive the spin of the BH or an upper limit to the radius of the NS.  \\citet{white86mnras} performed a systematic analysis of six NS LMXBs with EXOSAT and found line emission in 83\\% of the objects. The lines were broad with FWHM of 0.8-1.3~keV and had centroid energies of 6.6-6.9~keV. Further the line properties did not show obvious correlation with luminosity. This led \\citet{white86mnras} to conclude that the only plausible broadening mechanism was Comptonization in a cloud with a Thomson depth of a few and an electron temperature close to 7~keV. \\citet{hirano87pasj} analysed ten NS LMXBs observed with {\\it Tenma} and found line emission from six of them. The line energy had a weighted average of 6.66\\,$\\pm$0.05~keV with \\ews\\ in the range 20-60~eV. A significant line width of 0.55~keV was only obtained for one source, though a width of $\\sim$1~keV could not be rejected for the other cases. They attributed the line emission to recombination processes in the accretion disc corona and explained its width as Compoton scattering in the region with an optical depth of 1 and an average temperature $\\sim$1~keV. \\citet{asai00apjs} extended the analysis of the iron K emission lines to twenty NS LMXBs observed with ASCA. They detected significant iron lines from about half of the sources. The average properties of the lines were a line centre of 6.56~keV and a finite width of $\\sim$0.5~keV (FWHM) in six of the sources. They also found a large scatter in the \\ew\\ of the detected lines, ranging between 10 and 170~eV. They concluded that the iron K lines are likely produced through the radiative recombination of photoionised plasma and explained the line width as a combination of line blending, Doppler broadening and Compton scatterings.  The picture outlined by the analyses based on EXOSAT, {\\it Tenma} and ASCA data changed with recent analyses of observations of NS LMXB observations by \\xmm, \\chandra\\ and \\suzaku. A number of skewed iron lines have been claimed and their origin attributed to relativistic effects \\citep[e.g.,][]{serx1:bhattacharyya07apjl, 1808:papitto09aa, 1705:disalvo09mnras, gx340:dai09apjl} or to Compton scattering in an optically thick medium expanding, or converging, at relativistic velocities \\citep[e.g.,][]{laurent07apj, titarchuk09apj}. However only some of the NS LMXBs showed such asymmetric lines while others showed symmetric lines \\citep[e.g.,][]{cackett08apj}.  A systematic analysis of the iron~K lines is at this stage fundamental in order to establish why some sources exhibit skewed lines while others show symmetric lines or why lines are only present in some of the sources. Having an answer to these questions would help to clarify the currently controversial origin of the lines.  With this aim we performed a systematic analysis of all the NS LMXBs observed by \\xmm\\ since the beginning of the mission and publicly available up to the 30th of September 2009. We excluded from this sample observations of dipping or ADC sources, i.e. with known inclinations above 70\\deg, since their analyses at the Fe band are complicated due to strong absorption in the line of sight. We chose \\xmm\\ to perform this study for the following reasons: firstly, its high effective area both below and above the iron~K band allows a good determination of the shape and width of the iron line as well as the continuum. Secondly, \\xmm\\ has provided most of the detections and well determined profiles of broad iron lines. Finally, \\xmm, after 10 years of operation, has an ample archive of public data which can be used to determine class properties in the least unbiased way.  In total we analysed 26 observations of 16 sources from the \\xmm\\ archive. We restricted our analysis to sources observed in the EPIC~pn Timing Mode. This mode is specially suited for the observation of bright sources. Therefore, by selecting all the observations performed in this mode we obtained a sample of spectra with the best possible statistics.  In what follows we first present the properties of the selected sample and our analysis method. Special attention was payed to eliminate pile-up effects and to the treatment of background. Then, we performed spectral fitting accounting for the excess emission at the Fe~K energy band. We finally discuss the implications of our analysis regarding the characteristics of the lines and their possible origin. In an appendix we compare the characteristics of the lines in this work with previous analyses of the same observations and discuss the reason for the discrepancies whenever found. ",
        "conclusions": "\\label{sec:discussion}  We performed a systematic analysis of 26 \\xmm\\ observations corresponding to all the NS LMXBs observed with EPIC~pn Timing mode and publicly available up until the 30th of September 2009 to establish the characteristics of the Fe~K line emission in these objects.  In seven observations we did not detect the source significantly. For the remaining nineteen observations we extracted one spectrum per observation. We payed special attention to the effects of pile-up and background subtraction.  \\subsection{Fe line emission}  We detected Fe~\\ka\\ line emission in 80\\% of the observations where the source was significantly detected.  The energy of the line has values between 6.53 and 6.97~keV, consistent with highly ionised iron, from \\fettwo\\ to \\fetsix. The width has values between 0.17 and 1.15~keV and the \\ew\\ between 17 and 190~eV. Only in four cases out of the fifteen lines reported in this work did we find a width near or above 1~keV and an \\ew\\ above 100~eV. Interestingly, there is a gap in the distributions of $\\sigma$ and \\ew\\ between these high values and the rest of the sample. Lines with high values of $\\sigma$ or \\ew\\ have also associated large errors. \\citet{vaughan08mnras} interpreted a similar finding in the context of detection of narrow lines from AGN as the lines with large \\ews\\ being most likely false detections. In the cases studied here the lines with widths near 1~keV are also in the limit of detectability of 3$\\sigma$ and the large errors most likely point to an inappropriate modelling of the continuum for \\osix, \\seventeentw, and \\sixteen.  Considering the whole sample, the Fe line has a weighted average energy of 6.67\\,$\\pm$\\,0.02~keV, a width ($\\sigma$) of 0.33\\,$\\pm$\\,0.02~keV and an \\ew\\ of 42\\,$\\pm$\\,3~eV.  Recently, broad skewed Fe~K emission lines from the disc were reported from LMXBs containing a NS \\citep[e.g.][]{serx1:bhattacharyya07apjl, gx349:iaria09aa, gx340:dai09apjl}. All the \\xmm\\ observations for which a broad skewed Fe~line has been reported were included in our sample (see Sect.~\\ref{sec:properties}). However, in contrast to previous analyses, we {\\it do not} need to invoke relativistic effects to explain its width, which could be due to mechanisms such as Compton broadening. The lines are {\\it equally well} fitted with the relativistic {\\tt laor} component or with a simpler Gaussian component. Further, the profiles shown in Fig.~\\ref{fig:gallery} do not show an asymmetric shape, as expected if the lines are emitted close to the NS and shaped by relativistic effects. The line profile is instead symmetric, similarly to the one found in dipping sources \\citep{ionabs:diaz06aa}. The major difference between the analysis presented in this work and previous works is a careful treatment of pile-up effects, common in the observations of bright LMXBs, and a different modelling of the continuum in some cases, which has a strong effect in the \\ew\\ of the lines.  As shown by \\citet{brandt94mnras}, in a scenario where the iron line is generated at the inner disc, we would expect the centroid and the width of the iron line to decrease as the accretion rate diminishes and the disc recedes. The behaviour of the equivalent width is somewhat more complex: the line equivalent width is expected to first increase with \\xil\\ up to a certain value at which it starts to decrease (see Fig.~5b from \\citet{brandt94mnras}). This is due to the emission properties of iron atoms as a function of the ionisation parameter (Matt et al. 1994). When the most abundant ions are \\fetfour--\\fetsix, line emission is more intense than in the neutral case, due to the greater fluorescent yield and the smaller photoabsorption cross-section at the line energy. However, if \\xil\\ increases further, the ion distribution becomes dominated by fully stripped atoms and the line emission decreases. In the observations presented in this paper, we do not observe a correlation between the centroid and the width of the iron line with luminosity, in agreement with a similar systematic analysis based on ASCA observations \\citep{asai00apjs}. In contrast, the three cases for which a large width and equivalent width are measured occur at relatively low luminosities and are actually associated to large errors. This seems to be at odds with the behaviour predicted by \\citet{brandt94mnras} for lines generated at the disc and suggests instead that the physical condition of the line-emitting region is rather similar among the LMXBs.  An interpretation of the line in the context of disc emission would imply that the line is produced far from the neutron star in order to explain its symmetric shape. However, in such a case it is difficult to justify the variations of the line as the disc inner radius changes.  Alternatively, the origin of the line broadening could be Compton scattering in a corona. Based on a systematic analysis of twenty~NS LMXBs with ASCA, \\citet{asai00apjs} concluded that the iron lines were likely produced through the radiative recombination of photoionised plasma, being the line width a result of the combination of line blending, Doppler broadening and Compton scattering. Combining such effects, \\citet{kallman89apj} estimated that \\ews\\ of up to 100 eV were attainable from an ADC using standard parameters. This in agreement with the values obtained in this work.  Finally, the line could originate in a partially ionised wind as a result of illumination by the central source continuum photons and broadened by electron downscattering in the wind environment \\citep[e.g.,][]{laurent07apj, titarchuk09apj}. The line profiles obtained via such mechanism are again asymmetric, similar to the ones produced by relativistic effects. Since the line profiles obtained here are highly symmetric, we do not attribute the line origin to the outflow based on these observations.  \\subsection{Continuum emission}  The X-ray continua for 80\\% of the observations were well fitted by a model consisting of a blackbody and disc-blackbody components absorbed by neutral material.  For 15\\% of the observations, a model consisting of a blackbody and power-law components absorbed by neutral material was preferred to the combination of blackbody and disc-blackbody components.  For one observation a fit with three components, namely, blackbody, disc-blackbody and power-law was required in order to get an acceptable \\rchisq.  For the fits with blackbody and disc-blackbody components the temperature of the blackbody component had values between 1.5 and 2.8~keV, except for \\seventeentw, for which a temperature of 3.8\\,$^{+8.4}_{-1.1}$~keV was found. The temperature of the disc blackbody component had lower values between 0.6 and 1.3~keV, except for \\seventeentw, for which a temperature of 1.9\\,$\\pm$\\,0.3~keV was found. This is consistent with the interpretation of the blackbody component as the boundary layer between the neutron star and the inner disc, which shows smaller radii and higher temperatures. \\citet{church01aa}, based on an ASCA survey, found a broad dependence of the blackbody luminosity on the total luminosity not previously known for LMXBs in general, but while they found that the blackbody luminosity was falling more rapidly than the total luminosity as the mass accretion rate decreases, we find a very tight correlation where the blackbody luminosity is very near the value of 50\\% of total luminosity at all luminosities. This may be due to the different continuum model used, blackbody and cutoff power-law components \\citep{church01aa} versus blackbody and disc blackbody components (this work). Similarly to \\citet{church01aa}, we find that the major change in blackbody emission results primarily from changes in the emitting area, not the temperature. The independence of the temperature of this component on the global value of Eddington ratio lends support to the theoretical suggestion that the boundary layer is radiation pressure supported \\citep{inogamov99,revnivtsev06aa}.  For the sources fitted with blackbody and power-law components, the temperature of the blackbody component had values between 0.6 and 2.2~keV and the index of the power-law was between 1.8 and 2.6.  We detected an excess of emission at $\\sim$1~keV in {\\it all} the observations where we could use the spectrum down to 0.7~keV, except for the two accreting millisecond pulsars \\xtee\\ and \\sax. We modelled such emission with a Gaussian component centred between 1.0 and 1.15~keV in 12 observations and as an edge between 0.84 and 0.86~keV in 2 observations. Provided the feature has an astrophysical origin, its appearance always at similar energy points to line emission, since a soft component due to a e.g. blackbody emission is expected to change its temperature for different sources. Its energy is consistent with emission from highly ionised Fe, from \\fettwenty\\ to \\fetfour, and/or \\neten. We did not find a clear correlation between the width of the line emission at $\\sim$1~keV and the Fe~\\ka\\ band. However, the fact that the same feature is not observed in simultaneous EPIC and RGS exposures \\citep[e.g.][]{1323:boirin05aa} indicates that the feature is likely blended emission from various ions and additionally broadened by the same mechanism that the Fe~\\ka\\ line.  In this work, we show that effects of pile-up can significantly affect the width, and therefore the physical interpretation, of the iron line in NS X-ray binaries. Furthermore, parallel works on \\xmm\\ and \\suzaku\\ observations of the black hole X-ray binary GX~339-4 demonstrated that, taking into account pile-up effects, the breadth of the line was consistent with a truncation of the disc \\citep{gx339:done10mnras}, or with a black hole with moderate spin \\citep{gx339:yamada09apjl}, in contrast to previous claims \\citep[e.g.][]{gx339:miller06apj,gx339:miller08apjl}. Therefore, we urge caution in using piled-up data for detailed spectral analysis."
    },
    "1005/1005.1566_arXiv.txt": {
        "abstract": "{Studying the density profiles of galaxy groups offers an important insight into the formation and evolution of the structures in the universe, since galaxy groups bridge the gap between single galaxies and massive clusters.} {We aim to probe the mass of \\object{SL2S\\,J02140-0535}, a galaxy group at $z=0.44$ from the Strong Lensing Legacy Survey (SL2S), which has uncovered this new population of group-scale strong lenses. } {We combine strong lensing modeling and dynamical constraints. The strong lensing analysis is based on multi-band HST/ACS observations exhibiting strong lensing features that we have followed-up spectroscopically with VLT/FORS2. To constrain the scale radius of an NFW mass profile that cannot be constrained by strong lensing, we propose a new method by taking advantage of the large-scale dynamical information provided by VLT/FORS2 and KECK/LRIS spectroscopy of group members. } {In constrast to other authors, we show that the observed lensing features in \\object{SL2S\\,J02140-0535}  belong to different background sources: one at  $z$= 1.7 $\\pm$ 0.1 (photometric redshift) produces three images, while the other at  $z$ = 1.023 $\\pm$ 0.001 (spectroscopic redshift) has only a single image. Our unimodal NFW mass model reproduces these images very well. It is characterized by a concentration parameter $c_{200}$ = 6.0 $\\pm$ 0.6, which is slightly greater than the value expected from $\\Lambda$CDM simulations for a mass of M$_{200}$ $\\approx$ 1 $\\times$ 10$^{14}$ M$_{\\sun}$. The spectroscopic analysis of group members also reveals a unimodal structure that exhibits no evidence of merging. The position angle of the halo is $\\theta$ = 111.6 $\\pm$ 0.2, which agrees with the direction defined by the luminosity contours. We compare our dynamic mass estimate with an independent weak-lensing based mass estimate finding that both are consistent.  % } {Our combined lensing and dynamical analysis of \\object{SL2S\\,J02140-0535} demonstrates the importance of spectroscopic information in reliably identifying the lensing features. Our findings argue that the system is a relaxed, massive galaxy group where mass is traced by light. This work shows a potentially useful method for constraining large-scale properties inaccessible to strong lensing, such as the scale radius of the NFW profile. } ",
        "introduction": "In the context of structure formation and evolution, galaxy groups play a key role, spanning the regime between individual galaxies and galaxy clusters. Therefore, detailed studies of this intermediate regime of the mass spectrum are relevant. Galaxy groups have been mostly studied using X-ray and optical tracers, as well as numerical simulations \\citep[\\emph{e.g.}][and references therein]{group1,group2,group3,group4,group5,group6,mamongroup,group7,group8,group9,fabio,elena05,jespergroup}. Only recently, thanks to the unsurpassed combined depth, area, and image quality of the Canada France Hawaii Telescope Legacy Survey (CFHTLS)\\footnote{\\tt http://www.cfht.hawaii.edu/Science/CFHLS/}, a new population of lenses has been uncovered, with Einstein radii between $\\sim 3\\arcsec$ and $\\sim 8\\arcsec$, generated by galaxy-group-scale dark-matter haloes \\citep[this subsample of the CFHTLS forms the Strong Lensing Legacy Survey\\footnote{\\tt http://www-sl2s.iap.fr/} (SL2S),][]{sl2s}. This new population effectively bridges the gap between single galaxies and massive clusters, opening a new window of exploration in the mass spectrum. The first representative sample of these lenses was presented by \\citet{paperI}. Since then, a detailed investigation of some groups has been carried out: SL2S\\,J02176-0513 \\citep{hong09} and SL2S\\,J08544-0121 \\citep{my0854}.  In this paper, we present a detailed analysis of \\object{SL2S\\,J02140-0535} at $z=0.44$, combining the dynamics of the group members, and both the strong and weak gravitational lensing.  Some researchers have demonstrated the advantages of combining the different lensing regimes with the constraints that come from the stellar kinematics of the central brightest cluster galaxy (BCG). For example, \\citet{sand02,sand04} combined strong gravitational lensing with the velocity dispersion profile of the BCG and found that the models favored a shallower slopes in the NFW profiles.  \\citet{new09} following the above works and adding weak-lensing analysis, sampled the dark matter profile over three decades in radius in the lensing cluster Abell\\,611. Despite all these works, on group scales the efforts in that direction are just beginning. First, because strong lensing in groups is a relatively new field of research, and second, owing to observational limitations, i.e., there are fewer lensing arcs in groups and detecting weak lensing by groups of galaxies is challenging.  \\citet{mck10} probed the inner part of a group-scale halo using strong gravitational lensing and stellar kinematics, and  \\citet{tha10} showed that is possible to combine strong lensing and galaxy dynamics to characterize the mass distribution and the mass-to-light ratios of galaxy groups.   \\object{SL2S\\,J02140-0535} was previously studied by \\citet{alardalone} using a perturbative reconstruction of the gravitational lens. He found that the local shape of the potential and density of the lens, inferred from the perturbative solution, reveals the existence of an independent dark component that does not follow light. He argued that this particular shape of the dark halo is due to the merging of cold dark matter halos. Given these findings, the group properties represent a challenge, and to describe its mass distribution and perform a detailed study of the object is of significant worth and deserves attention. Our first aim in this work is to build a strong lensing mass model for \\object{SL2S\\,J02140-0535} using lensed features observed in HST/ACS data combined with ground-based spectroscopy. The second is to combine the strong lensing constraints with the dynamics of the groups members to study the mass profile out to the scale radius. Finally, we wish to compare the masses obtained with the three methods, namely, strong lensing, weak lensing, and group dynamics, to shed some light on \\object{SL2S\\,J02140-0535}.      The paper is organized as follows. In Sect.\\,\\ref{data}, we describe the observational data images and spectroscopy and present the multiple image identification. In Sect.\\,\\ref{dynamics}, we define how the dynamical constraints (on galaxy and group scales) enter into our strong lensing model. We also depict the profiles for both, the group and the galaxy-scale mass components. In Sect.\\,\\ref{model}, we present the strong lensing as well as the weak lensing analysis of \\object{SL2S\\,J02140-0535} and the main results of the work. In Sect.\\,\\ref{discuss}, we summarize and discuss our results. Finally in Sect.\\,\\ref{conclus}, we present the conclusions. All our results are scaled to a flat, $\\Lambda$CDM cosmology with $\\Omega_{\\rm{M}} = 0.3, \\Omega_\\Lambda = 0.7$ and a Hubble constant \\textsc{H}$_0 = 70$ km\\,s$^{-1}$ Mpc$^{-1}$. All images are aligned with WCS coordinates, i.e. north is up, east is left. Magnitudes are given in the AB system. ",
        "conclusions": "We have probed the gravitational potential of \\object{SL2S\\,J02140-0535} on a large radial range using complementary techniques: strong lensing, weak lensing, and dynamics. In the strong lensing regime, we performed a fit using an NFW profile and three galaxy-scale mass components as perturbations to the group potential. We used PIEMDs for the individual galaxies and set constraints on their geometrical and dynamical parameters using observational data (image and spectroscopy). We measured the velocity dispersion of the group using Keck (LRIS) and VLT (FORS2) spectra,  and we applied this information to constrain and  build a reliable strong lensing model based on one multiply imaged system and one single image system. In the weak lensing regime, we used the analysis reported in \\citet{paperI}. Our main results can be summarized as follows:   \\begin{enumerate}  \\item We have shown that the AB and C systems belong to different sources, one at $z$ = 1.7 $\\pm$ 0.1 (photometric redshift) and the other one at $z$ = 1.023 $\\pm$ 0.001 (spectroscopic), respectively.    \\item Our best-fit lens model reproduce quite well the image systems observed in the field of \\object{SL2S\\,J02140-0535}, with no extra assumptions about the underlying mass distribution (e.g. bimodal distribution). This agrees with the spectroscopic analysis that shows a well-defined group. Furthermore, we obtained a position angle of the halo of $\\theta=111.6\\pm0.2$, that is consistent with the orientation of the luminosity contours and the spatial distribution of confirmed members.  \\item The spectroscopic information provide constraints on the scale radius of the NFW profile. In the present work, we showed that is possible to use the dynamics at group scale to discriminate between a great range of possible $r_s$ values and $\\sigma$ (in particular  100 kpc $\\leq$ $r_s$ $\\leq$ 200 kpc and 523 km\\,s$^{-1}$ $\\leq$ $\\sigma$  $\\leq$ 737 km\\,s$^{-1}$, respectively).    \\item The scale radius and the velocity dispersion found in our best-fit models, $r_s$ = 170 $\\pm$ 18 kpc and $\\sigma$ = 664 $\\pm$ 18 km\\,s$^{-1}$, respectively, yields a  concentration value $c_{200}$ = 6.0 $\\pm$ 0.6, which is consistent with $c_{NFW}$ = 5 $\\pm$ 3, but slightly greater that the one predicted by $\\Lambda$CDM simulations.   \\item  The masses obtained with the strong lensing, weak lensing, and $M_{2D}$ methods agree within the errors, but the virial mass is slightly below these values. Thus, we demonstrated that it is possible to combine these  different methodologies to get a complete insight into the mass of a galaxy group.       \\end{enumerate}        The lensing and dynamical analysis of  \\object{SL2S\\,J02140-0535} presented in this paper,  has had two principal points. First, it has investigated the impact of including spectroscopic information of arcs (i.e. secure arc identifications) in strong lensing models to avoid complex mass distribution models that do not follow the light. Secondly, we studied the possibility of constraining large scale properties that could not be measured by strong lensing, such as the scale radius in the NFW profile. Future spectroscopic follow-up of the arcs A and B will provide confirmation of our model. In addition, the deepest HST images of the group will allow us to construct more accurate weak lensing models capable of probing an NFW profile on large scales, as well as  improving our strong lensing model."
    },
    "1005/1005.2942_arXiv.txt": {
        "abstract": "It has been proposed that primordial gas in early dark matter halos, with virial temperatures $T_{\\rm vir}\\gsim 10^4$K, can avoid fragmentation and undergo rapid collapse, possibly resulting in a supermassive black hole (SMBH).  This requires the gas to avoid cooling and to remain at temperatures near $T\\sim 10^4$ K.  We show that this condition can be satisfied in the presence of a sufficiently strong primordial magnetic field, which heats the collapsing gas via ambipolar diffusion.  If the field has a strength above $\\mid B\\mid\\gsim 3.6$ (comoving) nG, the collapsing gas is kept warm ($T\\sim 10^4$ K) until it reaches the critical density $n_{\\rm crit}\\approx10^3 {\\rm cm^{-3}}$ at which the roto--vibrational states of ${\\rm H_2}$ approach local thermodynamic equilibrium.  ${\\rm H_2}$--cooling then remains inefficient, and the gas temperature stays near $\\sim 10^4$K, even as it continues to collapse to higher densities. The critical magnetic field strength required to permanently suppress ${\\rm H_2}$--cooling is somewhat higher than upper limit of $\\sim 2$nG from the cosmic microwave background (CMB). However, it can be realized in the rare $\\gsim(2-3)\\sigma$ regions of the spatially fluctuating $B$--field; these regions contain a sufficient number of halos to account for the $z\\approx6$ quasar BHs. ",
        "introduction": "The discovery of very bright quasars, with luminosities $\\ge 10^{47}~{\\rm erg~s^{-1}}$, at redshift $z\\simeq 6$ in the Sloan Digital Sky Survey (SDSS) suggests that some SMBHs as massive as a few$\\times10^9~\\msun$ already existed when the universe was less than 1 Gyr old (see, e.g., Fan 2006 for a review).  The presence of these SMBHs presents a puzzle.  Metal--free stars, with masses $\\sim 100~\\msun$, are expected to form at redshifts as high as $z\\gsim 25$ (Abel et al. 2002; Bromm et al. 2002; Yoshida et al. 2008), and leave behind remnant BHs with similar masses (Heger et al. 2003).  However, the natural time-scale, i.e. the Eddington time, for growing these seed BHs by $\\gsim 7$ orders of magnitude in mass is comparable to the age of the universe (e.g. Haiman \\& Loeb 2001). This makes it difficult to reach $10^9~\\msun$ without a phase of rapid (super--Eddington) accretion, unless a list of optimistic assumptions are made in hierarchical merger models, in which multiple seed BHs are allowed to grow without interruption, and to combine into a single SMBH (Haiman 2004; Yoo \\& Miralda-Escud\\'{e} 2004; Bromley et al. 2004; Shapiro 2005, Volonteri \\& Rees 2006; Li et al. 2007; Tanaka \\& Haiman 2009).  An alternative class of explanations involves rapid gas accretion or collapse. In this family of models, primordial gas collapses rapidly into a SMBH as massive as $10^4-10^6~\\msun$ (Oh \\& Haiman 2002; Bromm \\& Loeb 2003; Koushiappas et al. 2004; Lodato \\& Natarajan 2006; Spaans \\& Silk 2006; Begelman et al. 2006; Volonteri et al. 2008; Wise \\& Abel 2008; Regan \\& Haehnelt 2009; Shang et al. 2010), possibly by accreting onto a pre--existing smaller seed BH (Volonteri \\& Rees 2005), or going through the intermediate state of a very massive star (Bromm \\& Loeb 2003), a dense stellar cluster (Omukai et al. 2008), or a ``quasistar'' (Begelman et al. 2008).  These so--called ``direct collapse'' models involve metal--free gas in relatively massive ($\\gsim 10^8~\\msun$) dark matter halos at redshift $z\\gsim 10$, with virial temperatures $T_{\\rm vir}\\gsim 10^4$K.  The gas that cools and collapses in these halos must avoid fragmentation, shed angular momentum efficiently, and collapse rapidly.  These conditions are unlikely to be met, unless the gas remains ``warm'', i.e. at temperatures $T_{\\rm vir}\\sim 10^4$K.  In particular, in recent numerical simulations, Shang et al. (2010) found that the gas in such halos, when collapsing in isolation, forms ${\\rm H_2}$ efficiently, and cools to temperatures of $T \\sim 300$ K.  Although no fragmentation was seen, the gas could ultimately fragment on smaller scales that have not yet been resolved (e.g. Turk et al. 2009).  More importantly, even if fragmentation was avoided, the cold gas flows inward at low velocities, near the sound speed of $\\sim 2-3~{\\rm km~s^{-1}}$, with a correspondingly low accretion rate of $\\sim 0.01~{\\rm M_\\odot~yr^{-1}}$. This results in conditions nearly identical to those in the cores of lower-mass minihalos; extensive ultra--high resolution simulations have concluded that the gas then forms a single $\\sim 100~{\\rm M_\\odot}$ star (Abel et al. 2002; Bromm et al. 2002; Yoshida et al. 2008) or perhaps a massive binary (Turk et al. 2009), rather than a supermassive star or BH.  ${\\rm H_2}$--cooling in early galaxies may be avoided when the gas is exposed to an intense UV flux $J$, either directly photo--dissociating ${\\rm H_2}$ (in the Lyman--Werner bands near a photon energy of $\\sim 12$eV) or photo--dissociating the intermediary ${\\rm H^-}$ (at photon energies $\\gsim 0.76$eV).  Requiring the photo-dissociation timescale, $t_{\\rm diss}\\propto J^{-1}$, to be shorter than the ${\\rm H_2}$--formation timescale, $t_{\\rm form}\\propto \\rho^{-1}$, generically yields a critical flux that increases linearly with density, $J^{\\rm crit} \\propto \\rho$.  Since the gas in halos with $T_{\\rm vir}\\gsim 10^4$K can cool via atomic Lyman $\\alpha$ radiation and loose pressure support, it inevitably collapses further. As a result, in these halos, the critical flux is high, $J^{\\rm crit}\\approx10^{2}-10^{5}$ (depending on the assumed spectral shape; Shang et al. 2010; see also Omukai 2001 and Bromm \\& Loeb 2003 who found similar, but somewhat higher values), which exceeds the expected level of the cosmic UV background at high redshifts.  Only a small subset of all $T_{\\rm vir}\\gsim 10^4$K halos, which have unusually close and bright neighbors, may see a sufficiently high flux (Dijkstra et al. 2008). In order to avoid fragmentation, the gas in these halos must also remain essentially free of any metals and dust (Omukai et al. 2008), which may be incompatible with the presence of such nearby luminous galaxies.  In this paper, we consider a different possibility to keep the gas warm, relying on heating by a primordial magnetic field.  Several mechanisms have been proposed to produce a global primordial magnetic field (PMF), with a field strength of order 1 (comoving) nG, during inflation and/or during various phase transitions in the early universe (see, e.g. Widrow 2002 and references therein).  If present, the PMF can be strongly amplified by flux--freezing inside a collapsing primordial gas, affecting ${\\rm H_2}$--formation and cooling. Sethi et al. (2008; hereafter S08) have shown that 0.2-2 nG fields can significantly enhance the ${\\rm H_2}$ fraction during the early stages of collapse. Schleicher et al. (2009) found similar results, and emphasized that at the high densities ($n\\gsim 10^8~{\\rm cm^{-3}}$) corresponding to later stages of the collapse, magnetic heating from 0.1-1nG fields results in significantly elevated temperatures.  In this paper, we consider field strengths higher than those in the two previous studies by S08 and Schleicher et al. (2009).  Our main result is that, in analogy with the UV--irradiation, there exists a critical magnetic field, which leads to a bifurcation of behaviors. If the PMF has a strength above $B>B_{\\rm crit}\\approx3.5$ nG, then the collapsing gas is kept warm ($T\\sim 10^4$ K) until it reaches the critical density $n_{\\rm crit}\\approx10^3 {\\rm cm^{-3}}$ at which the roto--vibrational states of ${\\rm H_2}$ approach local thermodynamic equilibrium.  ${\\rm H_2}$--cooling then remains inefficient, and the temperature stays near $\\sim 10^4$K, even as the gas collapses further.  On the other hand if $B<B_{\\rm crit}$, ${\\rm H_2}$--cooling is delayed, but the gas eventually cools down below $\\sim 1000$K.  The critical magnetic field strength we find is a factor of $\\sim$two higher than the existing upper limit from the CMB anisotropies (Yamazaki 2010; see more discussion in \\S~5 below).  However, it can be realized in the rare $\\gsim(2-3)\\sigma$ regions of the spatially fluctuating $B$--field. As we argue, the abundance of halos located in these high--field regions is sufficient to explain the number of the $z\\approx6$ quasars observed in the SDSS. ",
        "conclusions": "We found a plausible novel mechanism to form high-redshift SMBHs by direct gas collapse in early dark matter halos, aided by heating from the dissipation of a primordial magnetic field.  The model avoids many of the assumptions required in earlier models (such as an extremely high UV flux and the absence of ${\\rm H_2}$ and of other molecules and metals), but it does require a large primordial magnetic field, and relies on metal--free primordial gas.  We expect that, in general, {\\em any other} heating mechanism, which can compete with atomic HI cooling in the collapsing halo, down to a density of $n\\sim 10^{3}~{\\rm cm^{-3}}$, would produce the same effect as the $B$--field utilized here.  Interestingly, our model requires a magnetic field strength that is close to the existing observational upper limits.  We therefore expect that the values $B_0 \\simeq \\hbox{a few} \\times 10^{-9} \\, \\rm G$, required to form SMBHs, could be detectable by the future CMB experiments.  The upward fluctuations of such a strong primordial magnetic field might also leave their trace on cosmic structure formation: because of the high magnetic Jeans mass in these regions, the formation of dwarf--galaxy sized halos, with masses below $M \\simeq \\hbox{a few} \\times 10^{10} M_\\odot$, will be prevented.  We expect this can lead to further constraints on our SMBH--formation scenario, analogous to constraints on warm dark matter models (which produce a similar suppression of small-scale structures; Barkana, Haiman \\& Ostriker 2001)."
    },
    "1005/1005.0947_arXiv.txt": {
        "abstract": "{Intermediate mass protostars provide a bridge between low- and high-mass protostars. Furthermore, they are an important component of the UV interstellar radiation field. Despite their relevance, little is known about their formation process.} {We present a systematic study of the physical structure of five intermediate mass, candidate Class 0 protostars. Our two goals are to shed light on the first phase of intermediate mass star formation and to compare these protostars with low- and high-mass sources.} {We derived the dust and gas temperature and density profiles of the sample. We analysed all existing continuum data on each source and modelled the resulting SED with the 1D radiative transfer code DUSTY. The gas temperature was then predicted by means of a modified version of the code CHT96.} {We found that the density profiles of five out of six studied intermediate mass envelopes are consistent with the predictions of the ``inside-out'' collapse theory. We compared several physical parameters, like the power law index of the density profile, the size, the mass, the average density, the density at 1000 AU and the density at 10 K of the envelopes of low-, intermediate, and high-mass protostars. When considering these various physical parameters, the transition between the three groups appears smooth, suggesting that the formation processes and triggers do not substantially differ.} {} ",
        "introduction": "\\label{sec:introduction}  Stars are usually catalogued in three groups according to their mass: low-mass ($\\lesssim$2 M$_\\odot$), intermediate mass (IM) and high-mass ($\\gtsim$8 M$_\\odot$). The high-mass stars are usually found to form in tight clusters \\citep[e.g.][]{Hillenbrand1998}, including many low-mass stars as well. Low-mass stars on their own can be formed in isolation or in loose groups of a few objects per cubic parsec \\citep{Gomez1993}. While there is reasonable consensus on the overall process leading to the formation of isolated, low-mass stars \\citep[see e.g.][]{Shu1987}, the situation for  high-mass star formation is much more confusing. The major problem is that a central object with M $\\gtsim$ 8 M$_\\odot$ produces a radiative pressure stronger than the gravitational pressure, which should stop the accretion. Therefore stars with a mass greater than $\\sim$8 M$_\\odot$ cannot be formed in a simple-minded fully gravitationally driven process and other mechanisms need to be considered. Studying the differences in observed environments, several authors have proposed that high-mass stars are formed by the coalescence of lower-mass stars, while other authors favor ``monolithic'' formation \\citep[see for example the recent review by][]{Beuther2007}, where turbulence or winds/outflows can play a major role \\citep{Mckee2003,Krumholz2005}.   The differences between low- and high-mass protostars are not only restricted to the dynamics of the formation mechanism. Notable differences also exist when it comes to the chemistry. \\citet{Bottinelli2006} compared the abundances of O-bearing complex organic molecules (COMs) present in Hot Cores surrounding high-mass protostars and in Hot Corinos around low-mass protostars. They concluded that the abundances of the O-bearing COMs, compared to methanol or formaldehyde, are about one order of magnitude higher in Hot Corinos than Hot Cores. In addition, significant differences in the abundances of molecules constituting the grain mantles (H$_2$O, CO, CO$_2$,CH$_3$OH, OCN$^-$) were observed between low- and high-mass protostars \\citep{Boogert2004}. Finally, in contrast to high-mass star forming regions, low-mass protostars present an extremely enhanced molecular deuteration, the so-called Super-Deuteration \\citep[e.g.][]{Ceccarelli2007}, not observed in high-mass star forming regions.   The IM stars, namely stars whose mass lies in the 2 to 8 M$_\\odot$ range, are crucial in studies of star formation because they provide the link between low- and high-mass stars \\citep{DiFrancesco1997,Mannings1997,Mannings2000} and therefore can help to determine whether different formation processes are required. IM stars are also found in clusters \\citep[e.g.][]{Testi1998,Neri2007,Fuente2007}, with a smooth transition towards the low-mass star, loose cluster regime for star masses around 3.5 M$_\\odot$ \\citep{Testi1999}. These authors also conclude that IM stars mark the transition from the low-density aggregates of $\\lesssim$ 10 stars per cubic parsec of T Tauri stars to the dense clusters of $\\gtrsim$ 10$^3$ stars per cubic parsec associated with massive stars. Finally, IM stars are among the most dominant sources for the interstellar FUV field \\citep[e.g.][]{Habing1968,Gondhalekar1975}, which regulates the phases of the ISM in the Galaxy and, in turn, the overall Galaxy star formation process and history. In this context, the study of the formation of the IM stars is not only interesting in itself, but it can also greatly help the debate on the formation of high-mass stars. Indeed, due to their intermediate position, the study of IM protostars provides crucial information about the transition between the two formation regimes, as well as on the limits of the low-mass and high-mass formation scenarios.  Despite their importance, very little is known about the formation and first evolutionary stages of IM stars. Although some individual IM protostars have been studied in detail \\citep{Fuente2007,Neri2007,Hogerheijde1999,Schreyer2002}, the situation is so bad that to date we neither have a satisfying sample of Class 0 IM objects, namely objects representing the first phases of stellar formation where the protostar is embedded in its envelope and its luminosity is dominated by the accretion luminosity, nor a systematic study of their physical structure, as in the case for low-mass and high-mass protostars. On the contrary, several previous studies have considered samples of low- and high-mass protostars. Relevant similar studies to the present IM one are those reported by \\citet{Jorgensen2002} and \\citet{Shirley2002}, who targeted 7 and 11 low-mass Class 0 protostars respectively, and the systematic studies by \\citet{Mueller2002}, \\citet{Hatchell2003}, \\citet{Williams2005} and \\citet{Vandertak2000}, who targeted 31, 10, 32, and 9 high-mass protostars respectively. These authors derived the physical structure of the envelope for each protostar using continuum observations coupled with a radiative transfer code assuming a 1D spherical model.  Combined, these studies cover a luminosity range of 1 $<$ L$_{\\odot} <$ 22, and 1$\\times$10$^{3}$ $<$ L$_{\\odot} <$ 2$\\times$10$^{6}$ for low- and high-mass protostars respectively, leaving a significant gap in luminosity and  knowledge.   This paper aims to fill this gap, providing a systematic study of the physical structure of the envelopes of a sample of five IM protostars. In addition, this study will provide a valuable input for future work on the chemical structure of IM protostars \\citep[e.g.][ in prep]{Alonso-albi2009in_prep}. The present work follows the method presented in \\citet{Crimier2009a} to study the IM protostars OMC2-FIR4. Namely, the combination of sub-millimeter to infrared observations with radiative transfer codes are used to derive the physical structure of each envelope. We selected a sample of five representative candidate Class 0 IM protostars: CB3-mm, Cep E-mm, IC1396 N BIMA 2, NGC7129 FIRS 2, and Serpens FIRS 1. These sources were selected using the following criteria: \\begin{itemize} \\item the luminosity is in the 30 to 1000 L$_{\\odot}$ range (corresponding to a $\\sim$ 3 to 9 M$_{\\odot}$ of ZAMS stars); \\item the continuum maps at submillimetre and/or millimetre wavelengths are available; \\item the source is closer than 2.5 kpc to the Sun and isolated, namely no source of similar brightness in the submillimetre within an angular distance of $\\sim$100$''$. \\end{itemize} The coordinates, the distances and the bolometric luminosities of the selected sources are listed in Table \\ref{basic_param}.  \\begin{table*}[h] \\begin{tabular}{|l|lllll|} \\hline Source               &   CB3       & Cep E       & IC1396 N    & NGC7129     & Serpens     \\\\ & mm          & mm          & BIMA 2      & FIRS 2      & FIRS 1      \\\\  \\hline RA (2000)      & 00:28:42.1  & 23:03:12.7  & 21:40:41.8  & 21:43:01.5  & 18:29:49.8  \\\\ Dec (2000)          & 56:41:59.4  & 61:42:27.4  & 58:16:13.5  & 66:03:25.0  & 01:15:18.4  \\\\ Distance (pc)        & 2500   & 730 & 750 & 1250 & 230   \\\\ Luminosity (L$_{\\odot}$) & 1000   & 100 & 150 & 500 & 33     \\\\ \\hline \\end{tabular} \\caption{The coordinates, the distances and the bolometric luminosities of the sources of the sample.\\label{basic_param}} \\end{table*}   This article is organised as follows. Section \\ref{sec:General_method} describes the observations used and the method and tools employed to derive the dust and gas density and temperature profiles. Section \\ref{sec:Results} first summarises the general results obtained toward the IM protostar sample. Next, for each source, background information, data available, and the results obtained are presented. In Sect. \\ref{sec:Discussion}, the physical structure of the envelopes of IM protostars obtained in this study are compared with those obtained in previous studies of low- and high-mass protostar envelopes. In this section we also discuss the limitations of our modeling method. Finally, Sect. \\ref{sec:Conclusions} concludes the article. ",
        "conclusions": "\\label{sec:Conclusions}   We have derived the physical structure of the envelopes of five IM protostars, with luminosities between 30 to 1000 L$_\\odot$.  The envelope dust density and temperature profiles were determined by means of the 1D radiative transfer code DUSTY, using all continuum observations from the literature. The analysis assumed that the density profiles follow a single index power law and obtained self-consistently the temperature profile.  The  best-fit envelope models well reproduce the observations, namely the sub-millimetre radial brightness profiles and the SED between $\\sim$60 $\\mu$m and 1.3 mm, for each source. However, the model underestimates the 24$\\mu$m emission by several orders of magnitude. We ran test models to better understand what ingredient is missing and conclude that a ``simple'' cavity is not enough to reproduce the 24$\\mu$m observations. Apparently, the missing ingredient is a low-opacity escape route plus a warm dust contribution inside the envelope (circumstellar disc, warm outflow-excavated cavity...).  The gas density and temperature profiles were derived by assuming a constant dust-to-gas ratio and by computing the gas thermal balance at each point within the envelope. Because the gas equilibrium temperature strongly depends on the water abundance in the interiors of the envelopes, we also computed the expected water emission for each source.  We found that the gas and dust are thermally coupled across the envelope with differences less than 5 K in three out of five sources.  In IC1396 N BIMA 2 and NGC7129 FIRS 2 the gas is colder than the dust by at most 40 K, in a small region just where the icy mantles are predicted to sublimate. The predicted water line fluxes are consistent with the upper limits derived by the ISO observations.  One of the major goals of the present study was to ``use'' the IM protostars as a bridge between the low- and high-mass protostars with the hope that this will aid our understanding of the star formation process at either end. When comparing the characteristics derived by the modelling of the envelopes of low-, intermediate, and high-mass protostars, it appears that there is a smooth transition between the various groups. This suggests that {\\it there are basically no different triggers or processes between these mass regimes}. The power law index $\\alpha$ is similar in all three groups of objects. The majority of the sources have $\\alpha$ between 1.5 and 2. This is consistent with the theory of isothermal collapse from an initially singular isothermal sphere, the so-called inside-out expansion-wave collapse (Shu 1977).  Regardless of the mass group, a few sources have $\\alpha$ lower than 1.5, pointing perhaps to the collapse of an initially logotropic, virialised sphere (Lizano \\& Shu 1989, McLaughlin 96, and McLaughlin 97). Finally, the luminosity (mass) of the star depends on the size of the envelope, but does not depend on the density at a given temperature (for example at 10 K or 100 K).  \\bigskip  {\\it Acknowledgements.} We warmly thank Patrick Hennebelle for helpful discussions. One of us (N.Crimier) is supported by a fellowship of the Minist\\`ere de l'Enseignement Sup\\'erieur et de la Recherche. We acknowledge the financial support by PPF and the Agence Nationale pour la Recherche (ANR), France (contract ANR-08-BLAN-0225). The James Clerk Maxwell Telescope is operated by the Joint Astronomy Centre on behalf of the Science and Technology Facilities Council of the United Kingdom, the Netherlands Organisation for Scientific Research, and the National Research Council of Canada. This paper has been partially supported by  MICINN, within the programme CONSOLIDER INGENIO 2010, under grant  \"Molecular Astrophysics: The Herschel and Alma Era -- ASTROMOL\" ( ref.:  CSD2009-00038) Doug Johnstone is supported by a Natural Sciences and Engineering Research Council of Canada (NSERC) Discovery Grant. \\bigskip"
    },
    "1005/1005.2173_arXiv.txt": {
        "abstract": "Isophotal analysis of M87, using data from the Advanced Camera for Surveys, reveals a projected displacement of $6.8\\pm0.8$~pc ($\\sim0\\farcs1$) between the nuclear point source (presumed to be the location of the supermassive black hole, SMBH) and the photo-center of the galaxy. The displacement is along a position angle of $307\\pm17$\\degr\\ and is consistent with the jet axis. This suggests the active SMBH in M87 does not currently reside at the galaxy center of mass, but is displaced in the counter-jet direction. Possible explanations for the displacement include orbital motion of an SMBH binary, gravitational perturbations due to massive objects (e.g., globular clusters), acceleration by an asymmetric or intrinsically one-sided jet, and gravitational recoil resulting from the coalescence of an SMBH binary. The displacement direction favors the latter two mechanisms. However, jet asymmetry is only viable, at the observed accretion rate, for a jet age of $>$0.1~Gyr and if the galaxy restoring force is negligible. This could be the case in the low density core of M87. A moderate recoil $\\sim$1~Myr ago might explain the disturbed nature of the nuclear gas disk, could be aligned with the jet axis, and can produce the observed offset. Alternatively, the displacement could be due to residual oscillations resulting from a large recoil that occurred in the aftermath of a major merger $\\le$1~Gyr ago. ",
        "introduction": "It is generally assumed that supermassive black holes (SMBHs) reside at the centers of their host galaxies. However, SMBHs can be significantly displaced from their central locations by asymmetric forces during a merger, or by a second SMBH \\citep{2006MmSAI..77..733K}. In addition, if a binary SMBH forms and coalesces, anisotropic emission of gravitational waves can result in initial impulsive kick velocities of several thousand \\kms \\citep[e.g.,][]{2007arXiv0710.1338P}. Even if the kick velocity is small enough for the coalesced SMBH to remain in the galaxy, {\\it N}-body simulations have shown that the SMBH can oscillate within the bulge for $\\sim1$~Gyr before coming to rest \\citep{2008ApJ...678..780G}. Alternatively, if the galaxy contains a radio source, the SMBH may experience sustained acceleration due to intrinsic asymmetries in jet power \\citep[e.g.,][]{2007ARep...51...97T}.  The most direct way to find a displaced SMBH is to observe a spatial offset between the SMBH and the center of its host galaxy. This requires data at the highest possible spatial resolution. E/S0 galaxies that are minimally affected by extinction and contain an active galactic nucleus (AGN) provide good candidates; bulge isophotes determine the position of the galaxy center and the AGN (point source) determines the position of the SMBH. M87 is an ideal target for a displaced SMBH search. It is nearby, has a regular bulge, is relatively free of dust, hosts an AGN and jet \\citep[e.g.,][]{2001ApJ...551..206P}, and has been extensively observed by the Hubble Space Telescope ({\\it HST}). In this {\\it Letter}, we report the discovery of a $6.8\\pm0.8$~pc projected displacement between the center of M87, as defined by the galaxy isophotes, and the SMBH. ",
        "conclusions": "Four displacement scenarios have been considered; jets, binaries, perturbers, and gravitational recoil. Perturbers cannot produce the observed displacement amplitude, and the observed jet-offset alignment must be fortuitous in the binary case. On the other hand, jet acceleration and gravitational recoil provide natural explanations for the jet-offset alignment.  These mechanisms predict displacement velocities that are either $\\lap10\\kms$ or $\\gap100\\kms$. For example, jet acceleration of the SMBH produces an rms velocity $v_\\mathrm{rms}~\\approx~0\\kms$ in the presence of a restoring force, and $\\Delta{v}~\\approx~20\\kms~t_6^{-1}$ without. For SMBH-core oscillations induced by a large and early kick, the rms velocity of the SMBH is \\begin{equation} v_\\mathrm{rms}\\approx \\sigma \\frac{r_\\mathrm{rms}}{r_c} \\approx 12\\ \\mathrm{km\\ s}^{-1} \\sigma_{300} r_{500}^{-1} \\frac{r_\\mathrm{rms}}{20\\ \\mathrm{pc}}, \\end{equation} independent of the time since the kick or the kick amplitude. However, a velocity of several hundred \\kms\\ is possible if the SMBH is undergoing a large amplitude oscillation following a recent ($\\le 1$ Myr) kick. A similar velocity is expected in the binary SMBH scenario.  It may be possible to distinguish between the high and low velocity displacement mechanisms by comparing the relative recessional velocities of stars at the galaxy center and gas around the SMBH. Unfortunately, this will be difficult; the low central surface brightness will limit the accuracy with which velocities can be determined from stellar absorption lines. Alternatively, wide field VLBA astrometry, referenced against background quasars \\citep{2009AAS...21342009D}, may allow the proper motion of the nuclear point source to be determined. However, it will be difficult to distinguish between the cluster motion of M87 and the intrinsic kinematics of the SMBH.  The central gas disk \\citep{1994ApJ...435L..35H,1994ApJ...435L..27F} may provide alternative constraints on the displacement velocity. For example, a displaced SMBH retains gas whose Keplerian velocity exceeds the kick velocity \\citep{2006MNRAS.367.1746M}. \\cite{1994ApJ...435L..27F} estimate the disk-like structure in M87 to extend to $\\sim80$~pc. Assuming this represents the portion of the disk that remains bound to the SMBH, we can infer an initial kick velocity $\\sim 400\\kms$. This would produce strongly shocked regions in the disk, as it passes through the ambient gas, that may explain several observed properties. For example, while the gas in the central $\\pm0\\farcs4$ region of the disk is photo-ionized \\citep{1997ApJ...489..579M,1999ApJ...527..733S} and has a closely Keplerian velocity field \\citep{1994ApJ...435L..35H,1997ApJ...489..579M}, at larger radii non circular motions and morphological distortions are present \\citep{1999LNP...530..278F,2004AAS...20511002B}. In addition, gas in the outer parts of the disk is evidently excited by a 265\\kms shock \\citep{1997ApJ...490..202D}, rather than the central UV continuum source.  Two further consequences from a kick velocity $v_\\mathrm{k}\\approx400\\kms$ are noted. First, the shocked gas may be hot enough ($T\\sim~2\\times~10^6$\\,K) to produce X-ray emission. Second, at this velocity the observed displacement is reached after $\\sim~25000(\\Delta{r}/10\\mathrm{pc})$~yr, i.e., significantly less than the minimum jet age discussed earlier. If this minimum jet age corresponds to the time since the kick, the SMBH cannot be on its initial outward trajectory. \\cite{2008ApJ...678..780G} show for $v_k<50\\%$ of the escape velocity, the SMBH motion is rapidly damped and undergoes only one large amplitude oscillation on a time-scale of 1~Myr . Therefore, it is possible the SMBH is on a return trajectory."
    },
    "1005/1005.2663.txt": {
        "abstract": "We present a significant improvement over our previous calculations of the cosmic string contribution to cosmic microwave background (CMB) power spectra, with particular focus on sub-WMAP angular scales. These smaller scales are relevant for the now-operational Planck satellite and additional sub-orbital CMB projects that have even finer resolutions. We employ larger Abelian Higgs string simulations than before and we additionally model and extrapolate the statistical measures from our simulations to smaller length scales. We then use an efficient means of including the extrapolations into our Einstein-Boltzmann calculations in order to yield accurate results over the multipole range $2\\leq\\ell\\leq4000$. Our results suggest that power-law behaviour cuts in for $\\ell\\gtrsim3000$ in the case of the temperature power spectrum, which then allows cautious extrapolation to even smaller scales. We find that a string contribution to the temperature power spectrum making up $10\\%$ of power at $\\ell=10$ would be larger than the Silk-damped primary adiabatic contribution for $\\ell\\gtrsim3500$. Astrophysical contributions such as the Sunyaev-Zeldovich effect also become important at these scales and will reduce the sensitivity to strings, but these are potentially distinguishable by their frequency-dependence. ",
        "introduction": "\\label{sec:intro}  Observations of the cosmic microwave background (CMB) radiation and the large-scale distribution of galaxies indicate that cosmic structure was seeded in the very early stages of the universe \\cite{Nolta:2008ih,Percival:2009xn}, consistent with the inflationary paradigm. However, the datasets leave room for significant effects due to the presence of cosmic strings at subsequent times \\cite{Bevis:2007gh, Battye:2006pk, Fraisse:2006xc, Wyman:2005tu}. These strings (see Refs. \\cite{Vilenkin:1994book, Hindmarsh:1994re, Sakellariadou:2006qs, Copeland:2009ga} for reviews) are particularly important in that they are predicted by many physically motivated inflation models, including brane inflation models in the context of string theory \\cite{Sarangi:2002yt, Jones:2003da} but also models rooted in Grand Unified Theory (GUT) \\cite{Jeannerot:2003qv}. Such models also predict other types of cosmic defect, including textures and semilocal strings \\cite{Vachaspati:1991dz,Hindmarsh:1991jq,Achucarro:1999it,Achucarro:2007sp}, and their CMB signals are \\cite{Pen:1997ae,Durrer:1998rw,Bevis:2004wk,Cruz:2007pe,Urrestilla:2007sf} also of great interest.  In previous works we calculated cosmic string CMB temperature and polarization power spectra using field theory simulations of the Abelian Higgs model \\cite{Bevis:2006mj, Bevis:2007qz} and used the results to fit models with both string- and inflation-induced anisotropies to CMB data \\cite{Bevis:2007gh}. It was found that the data (which was primarily the WMAP third-year release \\cite{WMAP3-T=Hinshaw:2006ia}) favoured a model with a fractional string contribution to the temperature power spectrum at multipole moment $\\ell = 10$ of $\\fd = 0.09 \\pm 0.05$ with the spectral tilt of the inflation-induced primordial perturbations being $\\ns = 1.00 \\pm 0.03$. The latter was in contrast to the inflation-only result of $\\ns = 0.951^{+0.015}_{-0.019}$ \\cite{WMAP3-C=Spergel:2006hy} and showed, along with Refs. \\cite{Battye:2006pk,Urrestilla:2007sf}, that the inclusion of topological defects could readily allow $\\ns>1$ under those data, thanks to a parameter degeneracy found in Ref. \\cite{Bevis:2004wk}.  In the present article we present a significant improvement in the CMB power spectrum predictions from string simulations, with particular emphasis on small angular scales. While the numerical simulations of the Abelian Higgs model in Ref. \\cite{Bevis:2006mj} yielded CMB results covering multipoles $\\ell \\approx 2 \\rightarrow 1000$, the full range of angular scales relevant for WMAP data, they did not have the dynamic range to accurately investigate finer angular scales. Higher multipoles are now of great interest thanks to good coverage by sub-orbital CMB experiments \\cite{CBI=Sievers:2009ah, ACBAR=Reichardt:2008ay, QUAD=Friedman:2009dt, SPT=Lueker:2009rx,ACT=Fowler:2010cy}, plus the imminent arrival of full-sky Planck data \\cite{Planck} ($\\ell \\lesssim 2500$), and we have therefore focused on yielding accurate string predictions for $\\ell\\lesssim4000$.  Our ability to now include smaller scales stems partially from improvements in computational facilities and the carefully-chosen initial conditions that we now employ in our Abelian Higgs simulations. But more importantly, we have made developments in the method used to yield the CMB power spectra from the statistical measures of the energy-momentum distribution in our simulations: the unequal-time correlation functions (UETCs) \\cite{Hindmarsh:1993pu, Pen:1997ae}. We now model and extrapolate the measured UETCs to smaller scales in order to mimic results from larger simulations and have developed an efficient means of including UETC results at these extrapolated scales into our CMB calculations.  The need for the above extrapolation can be understood as follows. The width of a cosmic string is microscopic (perhaps at the GUT scale) while their separation at times of cosmological importance is of order the Hubble distance and hence it is not possible to simultaneously resolve both scales in numerical simulations during the epochs when strings would have impacted on the CMB. We can solve this problem by using \\emph{scaling} \\cite{Kibble:1984hp}: when the horizon is greater than about one hundred times the string width, we observe that the strings enter a late-time attractor solution in which statistical measures of the string distribution, when measured in horizon units, are constant in time. For example the average string length in a horizon volume divided by the horizon size is $\\approx50$. Similarly, the UETC measurements are independent of time once scaled by the horizon size and hence scaling provides knowledge of their values at times of cosmological importance. However, the measured UETC power is attenuated on scales close to the string width and scaling is broken for such length scales. Hence the resultant CMB results would show too little power at very small angular scales: the sharp changes in temperature caused by the Gott-Kaiser-Stebbins (GKS) effect \\cite{Gott:1984ef, Kaiser:1984iv,Stebbins:1987va} would have been smeared out by having been effectively sourced by strings whose width was around one-hundredth of the horizon size in the post-recombination era. For our previous work this was of little actual concern since those articles focused upon WMAP scales ($\\ell\\lesssim1000$). It is, however, necessary to extrapolate the UETC results to sub-string scales in order to obtain accurate results from field theoretic simulations for much higher multipoles.  Other approaches to calculating CMB perturbations from strings include using simulations of the Nambu-Goto (NG) type \\cite{Contaldi:1998mx, Landriau:2003xf, Fraisse:2007nu} and the unconnected segment model (USM) \\cite{Albrecht:1997nt, Pogosian:1999np, Battye:2006pk, Pogosian:2008am}. Nambu-Goto simulations can yield a greater dynamic range than field theoretic simulations but must still invoke scaling to yield CMB results over a wide range of scales. Furthermore the network correlation length used in the initial conditions appears to persist and may limit the reliable resolution of Nambu-Goto simulations \\cite{Ringeval:2005kr, Olum:2006ix}. The USM represents the string network as a stochastic set of moving ``sticks\", which disappear at an appropriate rate in order to give the chosen string scaling density. While it has no true dynamical content, it is computationally cheap and offers the flexibility to choose the coarse-grained network properties \\cite{Pogosian:1999np,Battye:2006pk,Pogosian:2007gi,Battye:2010hg,Battye:2010xz}.  The advantage of the Abelian Higgs model is that it includes the small-scale physics near the string width, which has a non-negligible impact on the string dynamics \\cite{Vincent:1997cx, Hindmarsh:2008dw}: energy from the strings is converted into massive gauge and Higgs radiation. In NG simulations this decay channel is not included and the string length density is significantly higher, with the long strings being converted into small loops that would then decay via gravitational radiation (although this process is not actually simulated). With the extra decay channel in the field theoretic simulations, decay via gravitational radiation would be less important, and this fact significantly changes the predictions for gravitational wave observations. In the case of the CMB on the other hand, it is the long strings that are important, and the key difference between the simulation results is the inter-string separation. A potential disadvantage of a field theory simulation is that computational constraints require the string width to be artificially increased in order to keep it above the simulation resolution, but we carefully show that this does not significantly affect the UETCs and therefore the CMB power spectra results.  In Sections \\ref{sec:method} and \\ref{sec:refine} we detail our methods, including an overview of the UETC approach, our field theory simulations, the tests of scaling that we employ and the sub-string extrapolation. We exhibit the resulting CMB power spectra in Section \\ref{sec:CMBresults}, and give conclusions in Section \\ref{sec:conclusions}.  %================================================================================ %================================================================================ %================================================================================ ",
        "conclusions": "\\label{sec:conclusions}  We have presented the first calculation of the cosmic string CMB temperature power spectrum contribution that is accurate over the multipole range $\\ell=2\\rightarrow4000$, a range that encloses the scales probed by the Planck satellite ($\\ell=2\\rightarrow2500$) and additionally scales at which sub-orbital data are becoming increasingly accurate. Further, we show that at $\\ell\\sim3000$ there is a knee in the power spectrum result due to the exponential decay of the early time contribution at small-scales, which then reveals the post-recombination component. This late-time contribution varies as roughly $1/\\ell$, which is the basic expectation for the GKS effect \\cite{Hindmarsh:1993pu}. Our results yield values for the power law exponent for each of the scalar, vector and tensor contributions to the overall high-$\\ell$ behaviour and can in principle be used to estimate the temperature contribution from strings for $\\ell>4000$. However this extrapolation must be performed with caution since our simulations do not have the dynamic range to yield these power laws with great confidence.  Our results indicate that the size of the cosmic string signal in the temperature power spectrum, relative the adiabatic inflationary component, is roughly 10 times greater at $\\ell\\approx3500$ than it is at $\\ell\\approx10$ or $\\ell\\approx400-1500$ and therefore small scales are particularly important for cosmic strings. While it is true that other contributions also become significant at small scales, for example the Sunyaev-Zel'dovich effect, such effects are frequency-dependent and hence can be identified and subtracted. Further, our results are not limited to temperature anisotropies but include polarization also, which is of great future importance for strings in the case of the B-mode. We present a comparison of our temperature and polarization results to the latest data in a separate, shorter article \\cite{Bevis:2010data}.  Our simulations also yield greater knowledge of the scaling properties of Abelian Higgs string networks, with improvements in the accuracy of results such as the scaling density and the UETCs. Additionally, we have measured the coherence function for Abelian Higgs strings for the first time, the form of which is important for calculations of this kind.  That our results stem from the Abelian Higgs model means, of course, that they are not necessarily accurate for cosmic superstrings (see eg. \\cite{Copeland:2009ga} for a review). These superstrings may have intercommutation probabilities significantly lower than that for gauge strings and additionally may form Y-shaped junctions, which do not form in the Abelian Higgs model for the parameters chosen here. While small-scale structure on the strings has been shown to lessen the impact of the intercommutation probability \\cite{Avgoustidis:2005nv}, the effect of Y-junctions on the network properties is highly uncertain. Both effects are likely to increase the string density and decrease the inter-string distance, resulting in a shift of the peak in the CMB signal to greater $\\ell$ values. This may therefore enable superstrings to be distinguished from conventional cosmic strings should a string component in the CMB be detected in future data.  Finally, we note that the power spectrum is not a complete statistical description of the anisotropies that would be seeded by cosmic strings, since there would be a significant non-Gaussian character created by the GKS effect. Higher order moments such as the bispectrum and trispectrum have been calculated \\cite{Hindmarsh:2009qk, Hindmarsh:2009es, Regan:2009hv}, in addition to realizations of CMB maps \\cite{Ringeval:2005kr, Landriau:2010cb}. While these calculations are challenging and either do not include recombination or are valid only for very limited multipole ranges, non-Gaussianity is an exciting channel for future cosmic string constraint or detection.  %================================================================================ %================================================================================ %================================================================================"
    },
    "1005/1005.2345_arXiv.txt": {
        "abstract": "In this paper, we use cubic polynomials to approximate RR Lyrae light curves and apply the method to $HST$ data of RR Lyrae stars in the halo of M31. We compare our results to the standard method of Fourier decomposition and find that the method of cubic polynomials eliminates virtually all ringing effects and does so with significantly fewer parameters than the Fourier technique. Further, for RRc stars the parameters in the fit are all physical in the sense that they can, in principle, be related to pulsation physics. Our study also reveals a number of additional periodicities in this data not found previously: we find 23 RR$c$ stars, 29 RR$ab$ stars and 3 multiperiodic stars. ",
        "introduction": "Even though a number of microlensing projects have yielded RR Lyrae light curves with excellent phase coverage \\citep[see, for example,][]{soz08}, phase coverage is more of a problem when studying RR Lyraes in external galaxies. In order to overcome this, a number of methods have been developed to use the observed data to get a smooth approximation to the actual RR Lyrae light curve that captures any real physical bumps or dips and eliminates any that are caused by numerics.  The traditional method applied has been Fourier decomposition wherein the observed data are fit by an expression of the form $$m(t) = A_0 + \\sum_{k=1}^{N}A_k\\sin (k{\\omega}t + {\\phi}_k). \\eqno(1)$$ Here $N$ is the order of the fit, $A_0$ is the mean magnitude, ${\\omega} = 2{\\pi}/P$, $P$ is the pulsating period and $t$ is the time of observation. Usually a least squares fitting procedure yields $A_0$, $A_k$, and ${\\phi}_k.$ A major problem with this technique is ringing: the Fourier curve given by equation (1) exhibits a series of unphysical bumps and dips. This can occur even when the original data are well distributed in phase but exhibit a large scatter. This is exacerbated when there are noticable gaps in phase coverage. The solution to this is to reduce the order of the fit but then this looses real features of the light curve.  Another method is that of Principal Component Analysis \\citep[PCA, see ][for examples]{kn04,deb09}. Here instead of sine functions being the basis, the data itself determines these basis functions: the resulting light curve is expressed as a sum of these basis functions. PCA has the advantage that a very good approximation of the light curve can be realised with significantly fewer parameters than required by the Fourier method. However, while a stringent comparison of the two methods is beyond the scope of this paper, we can say that, at least in the case of RR$c$ stars, the parameters of a fit using the methods developed here do have some physical interpretation.  \\citet{ak94} studied the method of cubic splines in approximating variable star light curves. A cubic spline is a series of cubic polynomials pieced together such that the intersection of two such polynomials is required to be continuous up to the second derivative. Our method uses cubic polynomials but does not require continuity up to the second derivative. Looking at figure 4 in \\citet{ak94}, the graphs for LCB12, LCR12, we again see ringing as a result of having 16 parameters: as our results show, this is a clear reason why cubics polynomials as opposed to cubic splines are to be preferred in approximating RR Lyrae light curves.  One motivation for this study is that requiring continuity up to the second derivative seems too stringent. Pulsation shocks are dramatic events with sudden reversals. There is no reason to suppose any fitted curve is continuously differentiable to the second degree. This is particularly true of fundamental mode RR Lyraes at phases close to minimum light where the star suddenly starts to get brighter. In this paper we examine the use of cubic polynomials to fit the light curves of RR Lyraes by requiring continuity only up to the first derivative. ",
        "conclusions": "We have found a new way to approximate the light curve of an RR Lyrae star by fitting cubic polynomials to the data. This method can fit the data with fewer parameters than Fourier series and suffers virtually no ringing. It can also estimate periodicities in the data. When we apply this method to RR Lyrae data in the Andromeda halo, we find, in addition to the multiperiodic star V90 reported by \\citet{br04}, an additional 2 other multiperiodic stars (V1 and V95, previously classified as type $RRc$) in the data sample: here we require this multiperiodicity to be present in both bands. Then the ratio of the number of RR$c$ stars to the ratio of the number of RR$abc$ stars is 0.442 - as opposed to \\citet{br04} who found a ratio of 0.462. In this ratio, \\citet{br04} do not count RRd stars in either numerator or denominator. The ratio of number of $RRc$ to number of $RRab$ where we include $RRd$ stars together with the $RRc$ stars is in this case is 0.473.    \\begin{table*} \\centering \\caption{Stars with multi-periodic components} \\label{NOLABEL} \\begin{tabular}{ccccccccccc}\\hline V  & Band & Period1 & $RA1$ & $SSE1$ & Period2 & $RA2$ & $SSE2$ & $pwr$ & $F$ & Number of points\\\\ \\hline  1  & F606W & 0.3815 & 0.856 & 0.252 & 0.5104 & 0.940 & 0.096 & 13.6 & 26.3 & 74  \\\\ 1  & F814W & 0.3816 & 0.779 & 0.227 & 0.5108 & 0.865 & 0.133 & 16.4 & 14.6 & 91 \\\\ 90 & F606W & 0.3533 & 0.823 & 0.579 & 0.4742 & 0.919 & 0.252 & 18.8 & 27.3 & 93\\\\ 90 & F814W & 0.3533 & 0.572 & 0.616 & 0.4747 & 0.809 & 0.263 & 26.2 & 30.2 & 99\\\\ 95 & F606W & 0.3616 & 0.781 & 0.429 & 0.4855 & 0.910 & 0.164 & 19.0 & 36.4 & 99\\\\ 95 & F814W & 0.3614 & 0.697 & 0.379 & 0.4855 & 0.744 & 0.309 &  9.7 & 6.0 & 115\\\\ \\hline \\end{tabular} \\end{table*}   \\begin{figure} \\centering \\epsfxsize=8.0cm{\\epsfbox{V1.eps}} \\caption{Results for multiperiodic star V1: an example of a multiperiodic star: the two dominant periods are $P_1 = 0.382$ and $P_0 = 0.510$ with a period ratio of 0.75.} \\label{UNIQUE_LABEL} \\end{figure}  \\begin{figure} \\centering \\epsfxsize=8.0cm{\\epsfbox{pc_minmax.eps}} \\caption{PC results at maximum (top) and minimum (bottom) light using a sixth order Fourier fit (open circles) and cubic polynomials (solid black squares).} \\label{UNIQUE_LABEL} \\end{figure}"
    },
    "1005/1005.5279_arXiv.txt": {
        "abstract": "{ In this lecture the basic theory of accretion disks is reviewed, with emphasis on aspects relevant for X-ray binaries and Cataclysmic Variables. The text gives a general introduction as well as a selective discussion of a number of more recent topics.} ",
        "introduction": "Accretion disks are inferred to exist as objects of very different scales: millions of km in low Mass X-ray Binaries (LMXB) and Cataclysmic Variables (CV), solar radius-to-AU scale disks in protostellar objects, and AU-to-parsec scale disks in Active Galactic Nuclei (AGN).  An interesting observational connection exists between accretion disks and jets (such as the spectacular jets from AGN and protostars), and outflows (the `CO-outflows' from protostars and the `broad-line-regions' in AGN).  Lacking direct (i.e. spatially resolved) observations of disks, theory has tried to provide models, with varying degrees of success. Uncertainty still exists with respect to some basic questions. In this situation, progress made by observations or modeling of a particular class of objects has direct impact for the understanding of other objects, including the enigmatic connection with jets.  In this lecture I concentrate on the more basic aspects of accretion disks, but an attempt is made to mention topics of current interest as well. Some emphasis is on those aspects of  accretion disk theory that connect to the observations of LMXB and CV's. For other reviews on the basics of accretion disks, see Pringle (1981), Papaloizou and Lin (1995).  For a more extensive introduction, the textbook by Frank et al. (2002). For a comprehensive text on CVs, Warner 1995. ",
        "conclusions": ""
    },
    "1005/1005.5423_arXiv.txt": {
        "abstract": "Helium core white dwarfs (WDs) with mass $M\\lesssim 0.20\\,M_{\\odot}$ undergo several Gyrs of stable hydrogen burning as they evolve.  We show that in a certain range of WD and hydrogen envelope masses, these WDs may exhibit $g$-mode pulsations similar to their passively cooling, more massive carbon/oxygen core counterparts, the ZZ~Cetis.  Our models with stably burning hydrogen envelopes on helium cores yield $g$-mode periods and period spacings longer than the canonical ZZ~Cetis by nearly a factor of 2.  We show that core composition and structure can be probed using seismology since the $g$-mode eigenfunctions predominantly reside in the helium core.  Though we have not carried out a fully nonadiabatic stability analysis, the scaling of the thermal time in the convective zone with surface gravity highlights several low-mass helium WDs that should be observed in search of pulsations:  NLTT~11748, SDSS~J0822+2753, and the companion to PSR~J1012+5307. Seismological studies of these He core WDs may prove especially fruitful, as their luminosity is related (via stable hydrogen burning) to the hydrogen envelope mass, which eliminates one model parameter. ",
        "introduction": "White dwarfs (WDs) are observed to pulsate in normal modes of oscillation ($g$-modes) which are determined by the structure of the stellar interior and atmosphere \\citep{win08}. Those with hydrogen atmospheres exhibit pulsations when they enter the ZZ Ceti variable (DAV) instability region, a discrete strip in the $T_{\\rm eff}$--$\\log g$ plane that spans $11,000\\,$K $ \\lesssim T_{\\rm eff} \\lesssim 12,250$\\,K at $\\log g \\approx 8.0$.  The ZZ~Ceti strip has been investigated both theoretically \\citep{bra97,wu99,fon03} and empirically \\citep{wes91,muk04,cas07,gia07}.  To date, all known ZZ~Ceti pulsators have masses $>0.5\\,M_\\odot$, implying cores composed of carbon, oxygen, and heavier elements.  Lower mass ($M<0.5\\,M_\\odot$) WDs with nearly pure helium cores are made on the red giant branch (RGB) when core growth is truncated before reaching $\\approx$0.45$-$0.47$M_{\\odot}$ ($\\log g \\approx 7.67$ at $T_{\\rm eff} \\approx 11,500$\\,K),  prior to the helium core flash  \\citep{dcr96,dom99,pie04}.  Two modes of envelope mass loss can cause this:  strong  winds or binary interaction.  Significant mass loss due to stellar winds in high metallicity systems may strip the H envelope,  preventing the He core flash \\citep{dcr96,han05}.  Common envelopes induced by binary interactions also lead to significant mass loss \\citep{ibe93,mar95}, and  make very low-mass He WDs ($M<0.2\\,M_{\\odot}$) when the binary interaction occurs at the base of the RGB \\citep{van96,cal98,bas06}.  Thus, He is the expected core composition for WDs below $\\approx$0.45$-$0.47$\\,M_{\\odot}$. However, very little direct evidence exists of He cores. The overbrightness of old WDs \\citep{han05} in the star cluster \\object[NGC 6791]{NGC~6791} \\citep{bed05} presents possible evidence.  The detection of low $\\log g$ young WDs make a plausible argument for the old WDs to be He core \\citep{kal07}, however, other possible explanations remain \\citep{del02,bed08a,bed08b,gar10}.  Asteroseismology offers the possibility of directly constraining the He core composition in these  low-mass WDs, as the $g$-mode periods  provide information on WD mass, mass of H envelope, and core composition \\citep{cor02,cas08}. Theoretically, these principles have been applied to C/O versus O/Ne core WDs by \\citet{cor04} and O versus He core WDs by \\citet{alt04}.  However, to carry this out, we need to find pulsating He core WDs, something that has yet to occur.  Past studies have illuminated a dichotomy in the evolution of the He core WDs \\citep{dri99,ser02,pan07} that impacts their seismic properties and prevalence as pulsators.  For masses $\\gtrsim$0.2\\,$M_{\\odot}$ (dependent upon metallicity) the H envelope experiences a multitude of H shell flashes that reduces its mass, eventually allowing the WD to cool rapidly. Such objects traverse the extrapolated ZZ~Ceti instability strip in $\\sim10$--$100$ Myr \\citep{pan07}, allowing for an investigation of their H layer mass, and confirmation of pure helium core. However, there are presently no known WDs in, or near, the extrapolated strip for masses in the $0.2\\,M_{\\odot}<M<0.5\\,M_{\\odot}$ range \\citep{ste08}, inhibiting such research. Less massive ($<0.2\\,M_\\odot$) helium WDs have a different evolution, undergoing stable H burning for Gyr, slowing their evolution to rates that may yield more in the extrapolated instability strip. However, the presence of a thick, actively burning hydrogen layer requires new seismic modeling, especially for the eventual assessment of the of the instability strip for these unusual WDs.  The recent discovery of three such objects \\citep{kaw09,kil09} makes our work quite timely.  Motivated by a desire to study the pulsational properties of these long-lived systems, we begin in Section\\ref{sec:bmod} by constructing He core WD models with stable H burning shells flexible enough for seismic investigations, and compare to the results from evolutionary codes \\citep{ser02, pan07}. We discuss the unusual seismic properties of these objects in Section \\ref{sec:panal}, where we calculate their adiabatic mode structure and, using an approximation for the instability criterion of Brickhill's theory \\citep{bri91}, highlight the potential  location of the He core WD instability strip. In Section \\ref{sec:conc}, we suggest a few intriguing pulsation candidates amongst the very lowest mass WDs \\citep{van96,cal98,bas06,kaw06,kaw09,kil09}, where our simple models apply. We close by highlighting the need for future work, especially if observations of our suggested targets yield the first pulsating, low-mass, He WD. ",
        "conclusions": "\\label{sec:conc}  Our work highlights the $\\log g$--$T_{\\rm eff}$ parameter space where observable pulsations may be present.  This reveals at least three pulsation candidates, \\object[NLTT 11748]{NLTT~11748} \\citep{kaw09,ste10}, \\object[SDSS J082212.57+275307.4]{SDSS~J0822+2753} \\citep{kil09}, and \\object[PSR J1012+5307]{PSR~J1012+5307} \\citep{van96,cal98}, all of which should be observed for variability on timescales of 200--1000\\,s.   Mode detections and a measurement of the mode period spacing would provide key evidence for an He core composition and large radius, as we predict a period spacing of $\\approx 90$\\,s for $\\ell = 1$, whereas in normal ZZ~Cetis, this number is $\\approx 50$\\,s \\citep{kle98,kan05,pec06}.  Our results regarding the region of instability should be confirmed through future nonadiabatic stability analyses, an issue that is beyond the scope of the present paper.  NLTT~11748 is  highlighted in Figure \\ref{fig:inststrip} as a candidate for observable pulsations.  Recent observations by \\citet{ste10} did not find pulsations down to a 5\\,mmag level, however, they did discover it to be the first eclipsing He WD system.  Given its measured $\\log g = 6.20$ and $T_{\\rm eff} = 8540$\\,K, our models predict a total mass of $0.17\\,M_{\\odot}$ and envelope mass of $3.15 \\times10^{-3}\\,M_{\\odot}$, comparable to that reported by \\citet{kaw09}.  Our numerical pulsation analysis reveals that for this object the lowest order $g$-mode ($\\ell = 1$ and $n = 1$) has a period of 245\\,s but more importantly the mean period spacing is 89\\,s for the $\\ell = 1$ modes and 51\\,s for the $\\ell = 2$ modes.  Figure \\ref{fig:propdiag} clearly shows these modes to preferentially reside in the core, offering a unique opportunity to probe the core composition of an He WD.  The current candidates were found in surveys that target other phenomena: the Sloan Digital Sky Survey (SDSS; \\citealt{kil07}), high velocity stars \\citep{kaw06,kaw09}, and companions to pulsars \\citep{van96,cal98,bas06}.  However, the survey selection criteria (photometric colors), although incomplete, can favor WDs of higher $T_{\\rm eff}$ ($T_{\\rm eff} \\gtrsim 11,000$\\,K in the SDSS for $\\log g < 6$, \\citealp{kil07}).  Surveys able to select low-gravity WDs down to $T_{\\rm eff} = 8,000$\\,K will significantly impact the study of pulsating He core WDs."
    },
    "1005/1005.2729_arXiv.txt": {
        "abstract": "{We report on the detection with the HIFI instrument on board the Herschel satellite of the two hydrogen chloride isotopologues, H$^{35}$Cl and H$^{37}$Cl, towards the massive star-forming region W3\\,A. The $J$=1--0 line of both species was observed with receiver \\textit{1b} of the HIFI instrument at $\\sim$625.9 and $\\sim$624.9\\,GHz. The different hyperfine components were resolved. The observations were modeled with a non-local, non-LTE radiative transfer model that includes hyperfine line overlap and radiative pumping by dust. Both effects are found to play an important role in the emerging intensity from the different hyperfine components. The inferred H$^{35}$Cl column density (a few times $\\sim$10$^{14}$\\,cm$^{-2}$), and fractional abundance relative to H nuclei ($\\sim$7.5$\\times$10$^{-10}$), supports an upper limit to the gas phase chlorine depletion of $\\approx$200. Our best-fit model estimate of the H$^{35}$Cl/H$^{37}$Cl abundance ratio is $\\approx$2.1$\\pm$0.5, slightly lower, but still compatible with the solar isotopic abundance ratio ($\\approx$3.1). Since both species were observed simultaneously, this is the first accurate estimation of the [$^{35}$Cl]/[$^{37}$Cl] isotopic ratio in molecular clouds. Our models indicate that even for large line opacities and possible hyperfine intensity anomalies, the H$^{35}$Cl and H$^{37}$Cl $J$=1-0 integrated line-intensity ratio provides  a  good estimate of the $^{35}$Cl/$^{37}$Cl  isotopic abundance ratio.} ",
        "introduction": "\\label{sec:introduction}  Chlorine has two stable isotopes ($^{35}$Cl and $^{37}$Cl) and an ionization potential of 12.97 eV (\\textit{i.e.,} slightly below that of hydrogen). Hence, it can be ionized by UV photons (912-956 \\AA) in diffuse clouds and in the edges of photon-dissociation regions (PDRs). Once ionized, Cl$^+$ reacts with molecular hydrogen exothermically to form HCl$^+$, a process that initiates the chemical reactions of chlorine. In cloud interiors, HCl$^+$ can be formed by reactions starting with neutral Cl and H$_3^+$. The chemistry of chlorine in interstellar clouds has been the subject of various works (see Neufeld \\& Wolfire 2009). These studies predict that hydrogen chloride (HCl) is the most abundant Cl-bearing molecule in dense clouds.   The HCl hyperfine lines (see Sect.~3.2) can be resolved in interstellar sources only with heterodyne receivers equipped with high spectral resolution spectrometers. The HCl $J$=1-0 line at $\\sim$625.9\\,GHz was first detected with the \\textit{Kuiper Airbone Observatory} towards Orion \\citep{Blake1985} and followed by detections in Sgr~B2 \\citep{Zmuidzinas1995}, several positions in OMC-1 \\citep{Schilke1995}, and in Mon~R2   using the \\textit{Caltech Submillimeter Observatory} with good atmospheric transparency (Salez et al. 1996; SFL96 hereafter). The inferred HCl column densities are in the range 10$^{13}$-10$^{14}$\\,cm$^{-2}$. SFL96 also presented the first detection of  H$^{37}$Cl  towards Orion. Since the chemical reactions involving HCl are relatively well understood (see Sect.~\\ref{chemistry}), the H$^{35}$Cl/H$^{37}$Cl abundance ratio should provide a good measure of the $^{35}$Cl/$^{37}$Cl isotopic ratio. Both $^{35}$Cl and $^{37}$Cl nuclei are believed to  form in the last burning stages of massive stars ($>$10\\,\\Msol) and by means of ``explosive nucleosynthesis\" during supernovae detonation (\\textit{e.g.,} Woosley \\& Weaver 1995). Therefore, observations of H$^{35}$Cl and H$^{37}$Cl, and accurate measurements of the $^{35}$Cl/$^{37}$Cl ratio in different environments, can provide some insight into the chemical evolution of both isotopes, thus into Galactic chemical evolution. The species HCl was also detected recently towards the carbon-rich evolved star IRC+10216 by \\citet{Cernicharo2010}.   Using HIFI, the \\textit{Heterodyne Instrument for the Far-Infrared} \\citep{deGraauw2010}, on board the \\textit{Herschel Space Observatory} \\citep{Pilbratt2010}, we present in this \\textit{Letter} the detection of the $J=1-0$ rotational transition of H$^{35}$Cl and H$^{37}$Cl towards the massive star-forming region W3. The broad frequency coverage of \\emph{HIFI} allows us to observe both isotopologues of HCl with the same relative calibration in a wide variety of astronomical environments. Here we present the first accurate determination of the $^{35}$Cl/$^{37}$Cl isotopic ratio by a detailed model of the excitation of the hyperfine levels developed using a non-local radiative transfer code. The effect of line overlaps between the hyperfine components, and radiative pumping by dust photons are discussed and modeled in detail. ",
        "conclusions": "\\label{sec:conclusions}   We have presented the first detection of H$^{35}$Cl and H$^{37}$Cl towards the W3\\,A H{\\sc ii} region. The inferred H$^{35}$Cl column density (a few times $\\sim$10$^{14}$\\,cm$^{-2}$) and fractional abundance ($\\sim$7.5$\\times$10$^{-10}$ per H nucleus) provide an upper limit to the gas phase chlorine depletion of $\\approx$200. This value is lower than that observed towards Orion hot core, but similar to that inferred towards Mon~R2 (SFL96). Radiative transfer models including HFS line overlap and pumping by dust photons have been used to interpret the observations. The best-fit model provides a H$^{35}$Cl/H$^{37}$Cl abundance ratio of $\\approx$2.1, which is both lower than the solar value ($\\approx$3.1) and lower than the previous estimate towards Orion ($\\approx$4-6; SFL96). On the other hand, it is similar to the [$^{35}$Cl]/[$^{37}$Cl] ratio obtained in the IRC+10216 circumstellar envelope  from [Na$^{35}$Cl]/[Na$^{37}$Cl] and [Al$^{35}$Cl]/[Al$^{37}$Cl] measurements \\citep{Cernicharo1987, Cernicharo2000}."
    },
    "1005/1005.3944_arXiv.txt": {
        "abstract": "We present an analysis of high-resolution spectra of six core-helium-burning `clump' stars in the open cluster NGC~6134. Atmospheric parameters ($T_{\\rm eff}$,  log~$g$,  $ v_{\\rm t}$, and [Fe/H]) were determined in our previous  study by Carretta et al.\\ (2004). In this study we present abundances of  C, N, O and up to 24 other chemical elements.  Abundances of carbon were  derived using the ${\\rm C}_2$ Swan (0,1) band head at 5635.5~{\\AA} (FEROS spectra) and the ${\\rm C}_2$  Swan (1,0) band head at 4737~{\\AA} (UVES spectra). The wavelength interval 7980--8130~{\\AA}, with strong CN features, was analysed in order to determine  nitrogen abundances and $^{12}{\\rm C}/^{13}{\\rm C}$  isotope ratios.  The oxygen abundances were determined from the [O\\,{\\sc i}] line at 6300~{\\AA}.  Compared with the Sun and other dwarf stars of the Galactic disk, mean abundances in the investigated clump stars  suggest that carbon is depleted by about 0.2~dex, nitrogen is overabundant by about  0.3~dex and oxygen is underabundant by about 0.1~dex. This has the effect of lowering the mean C/N ratio to $1.2\\pm0.2$.  The mean $^{12}{\\rm C}/^{13}{\\rm C}$ ratios are lowered  to  $9\\pm2.5$.   Concerning other chemical elements, the analysis of sodium and magnesium lines (in NLTE), lines of other $\\alpha$-elements, iron-group and heavier chemical elements gave abundance ratios close to the solar ones. ",
        "introduction": "Open clusters are important tools for the study of the Galactic disk as well as for understanding stellar evolution (e.g.,  Friel et al. 2002; Bragaglia et al. 2008;  Jacobson et al. 2009; Santos et al. 2009). They are the best tool to understand whether and  how the slope of the radial metallicity distribution changes with time, since they have formed at all epochs  and their ages, distances and metallicities are more accurately derived than for the  field stars. Open clusters are excellent laboratories for investigations of  stellar evolution as well. Since cluster members were initially of approximately identical chemical composition, all changes in stellar atmospheres of evolved stars  are related to internal and external processes of stellar evolution (see, e.g., Pallavicini 2003 and references therein). Changes of the abundances of carbon, nitrogen and oxygen are most often seen in  evolved stars.  The enhancement of CN bands and altered carbon isotope ratios in evolved stars of open clusters were already reported 30 years ago (e.g. Pagel 1974;  McClure 1974). However, the detailed analyses of abundances in stars of open  clusters from high-resolution spectra are still necessary for understanding the  processes of dredge-up and extra-mixing affecting the chemical composition of  atmospheres in evolved low-mass stars. Detailed spectral analyses of CNO elements in  stars of open clusters are still rather scarce (Gilroy 1989; Gilroy \\& Brown 1991;  Luck 1994; Gonzalez \\& Wallerstein 2000; Tautvai\\v{s}ien\\.{e} et al.\\ 2000, 2005;  Origlia et al. 2006; Gratton et al.\\ 2006; Smiljanic et al.\\ 2009, etc.).  The Bologna Open Cluster Chemical Evolution project (BOCCE) is dedicated to constraining the formation and chemical evolution of the Galactic disk (Bragaglia \\& Tosi 2006, Carretta et al.\\ 2007, and references therein).  This paper is a part of the effort to determine detailed elemental abundances in open clusters, with the final goal of deriving the time evolution of abundances in the Galactic disk. In particular, we concentrate here on carbon and nitrogen, but also derive abundances of more than 20 other chemical elements.  The open cluster NGC~6134 is an intermediate-age, moderately concentrated open cluster (Trumpler class II3m) located almost on the galactic plane ($\\alpha_{2000}=16^{h}27.8^{m}, \\delta_{2000}=-49^{\\circ}09.4^{\\prime}.9; l = 324.91^{\\circ}, b = -0.20^{\\circ}$). The first extensive photometric study was published by Lindoff (1972) who derived a colour excess $E(B-V)=0.45$, a distance of about 700~pc, and an age of about 0.7~Gyr. These values were based only on $UBV$ photographic data. Kjeldsen \\& Frandsen (1991) published $UBV$ CCD data for 66 stars at the center of the cluster and obtained $E(B-V)=0.46$ and $m-M=11.25$. Coravel radial velocity measurements and photometry in the $UBV$ and $CMT_{1}T_{2}$ system of 24 red giants, supplemented by $DDO$ observations of 11 stars, were carried out by Claria \\& Mermilliod (1992) for membership and binarity analysis, who identified 17 red giant members and 6 spectroscopic binaries. The mean cluster radial velocity was found to be $-26.0\\pm 0.24$~km~s$^{-1}$, the reddening  $E(B-V)=0.35\\pm 0.02$, and the distance about 760~pc. The weighted mean value of ${\\rm [Fe/H]}=-0.05\\pm 0.12$ was evaluated from the UV excesses. Str\\\"{o}mgren photometry was analysed by Bruntt et al.\\ (1999). They determined $E(b-y)=0.263\\pm 0.004$ ($E(B-V)=0.365$), [Fe/H]$=0.28\\pm 0.02$, and age=$0.69\\pm 0.10$~Gyr. The colour-magnitude diagram (Bruntt et al.\\ 1999) shows a \"clump\" of core-He-burning stars, several red giant branch (RGB) stars, and a main sequence. From $BVRI$ CCD observations Ahumada (2002) has determined a colour excess $0.29 < E(B-V) < 0.37$, age of 1.26~Gyr, and a distance of about $1080\\pm 50$~pc, which is larger than in the previous analyses.  \\input epsf \\begin{figure} \\epsfxsize=\\hsize \\epsfbox[0 0 500 430]{Mik-fig1.eps} \\caption{Field of $10\\times10$~arcmin$^2$ centered on NGC 6134, with the programme stars indicated by their numbers according to Lindoff (1972).} \\label{fig1} \\end{figure}   \\input epsf \\begin{figure} \\epsfxsize=\\hsize \\epsfbox[-5 -10 600 550]{Mik-fig2.eps} \\caption{The colour-magnitude diagram of the open cluster NGC~6134. The core-He-burning `clump' stars analysed are indicated by the filled squares. The diagram is based on Str\\\"{o}mgren photometry by Bruntt et al.\\ (1999)} \\label{fig1} \\end{figure}   Precise iron abundances from high resolution spectra for six stars in this cluster have been determined by Carretta et al.\\ (2004). An overall metallicity ${\\rm [Fe/H]}=0.15\\pm 0.03$ with $rms=0.07$ was found.  Carretta et al. also computed the reddening from the temperatures -derived spectroscopically- and the Alonso et al. (1999) colour-temperature relations, finding $E(B-V)=0.363\\pm0.014$, in very good agreement with Bruntt et al.\\ (1999). Using the same spectra and method of analysis, in this paper we continue detailed abundance investigations for the clump stars NGC 6134\\_39, 69, 75, 114, 129, and 157. Elemental abundances in three other clump stars of this cluster (NGC 6134\\_30, 99, and 202) have been investigated recently by Smiljanic et al.\\ (2009). ",
        "conclusions": "The abundances relative to hydrogen [El/H]\\footnote{In this paper we use the customary spectroscopic notation [X/Y]$\\equiv \\log_{10}(N_{\\rm X}/N_{\\rm Y})_{\\rm star} - \\log_{10}(N_{\\rm X}/N_{\\rm Y})_\\odot$} and $\\sigma$ (the line-to-line scatter) derived for up to 26 neutral and ionized species for the programme stars are listed in Table~3. The average cluster abundances and dispersions about the mean values for NGC\\,6134 are presented in Table~3 as well.  \\subsection{Galactic radial abundance gradient}  The open cluster radial abundance gradient was analysed and discussed many times during several decades (see e.g. Friel 1995, Twarog et al.\\ 1997; Bragaglia et al.\\ 2001, 2008; Friel et al.\\ 2002, 2005; Carretta et al.\\ 2004, 2005, 2007; Yong et al.\\ 2005; Sestito et al.\\ 2006, 2007, 2008; Jacobson et al.\\ 2008, 2009; Smiljanic et al.\\ 2009; Pancino et al.\\ 2010 and references therein). Twarog et al.\\ (1997) first proposed that the open cluster abundance distribution is not a negative linear gradient but two separate distributions, each of constant metallicity, divided at $R_{\\rm gc}= 10$ kpc. Clusters in the inner part are of solar metallicity, while in the outer part the mean metallicity is about $-0.3$~dex. The recent investigations of open clusters, reaching also the more distant parts of the disk, show that maybe the gradient is not the same in all the disk: it is rather steep in the inner disk ($R_{\\rm gc}< 12-14$~kpc), and it flattens in the outer disk (c.f., Carraro et al.\\ 2004, 2007; Yong et al.\\ 2005; Sestito et al.\\ 2006, 2008; Jacobson et al.\\ 2009).  \\input epsf \\begin{figure*} \\epsfxsize=\\hsize \\epsfbox[0 0 600 360]{Mik-fig4.eps} \\caption {Radial distribution of open cluster abundances. Open squares show results by Bragaglia et al.\\ (2001, 2008), Carretta et al.\\ (2004, 2005, 2007), and Sestito et al.\\ (2006, 2007, 2008). The result of this study is marked as a filled square. Crosses show results presented in the papers by Friel et al.\\ (2005) and Jacobson et al.\\ (2008, 2009), pluses -- Yong et al.\\ (2005), triangles -- Smiljanic et al.\\ (2009), and reversed triangles -- Pancino et al.\\ (2010). } \\label{rad.distr.} \\end{figure*}  Recently, Magrini et al.\\ (2009)  used a set of literature abundances of open clusters based on high resolution spectroscopy to compare the gradient, and its time evolution, with their chemical evolution models.  However, their sample is not homogeneous (distances, ages, and abundances were taken from papers of many different groups) so it is not an ideal set, since systematics can mask or produce features in the distribution. The BOCCE project aims at collecting a homogeneous sample, well spread over the Galactic plane, covering all ages, and metallicities. Fig.~4 displays the radial distribution of some elemental abundances for BOCCE clusters analysed so far, and for others in recent studies. The scatter is quite large at all radii, but NGC\\,6134 agrees well with results of other open clusters at the same $R_{\\rm gc}$ of 7.5 kpc.   \\subsection{Carbon, nitrogen and oxygen}  Since the [C\\,{\\sc i}] 8727~{\\AA} line is blended with CN in spectra of red giants, and since the probability of increased  strengths of CN molecular lines is present, the ${\\rm C}_2$ Swan (0,1) band head  at 5635.5~{\\AA} is a quite popular feature for the carbon abundance determinations. It  was used in our previous studies of giants (e.g. Tautvai\\v{s}ien\\.{e} et al.\\ 2000, 2005). In Fig.~5, a fit to the NGC\\,6134\\_157 spectrum at ${\\rm C}_2$ 5635.5~{\\AA} is shown.  The Swan band (1,0) with the head at 4737~{\\AA} is rarely used for carbon abundance determinations,  since the region is quite crowded by additional lines. Fortunately, two features of ${\\rm C}_2$ at  4732.8~{\\AA} and 4735.3~{\\AA} were clearly seen in the spectra obtained with the UVES spectrograph.  In Fig.~6, we show an example of spectral fit to these ${\\rm C}_2$ bands for the star NGC\\,6134\\_75.  \\input epsf \\begin{figure} \\epsfxsize=\\hsize \\epsfbox[-20 -20 620 470]{Mik-fig5.eps} \\caption {Small region of NGC\\,6134\\_157 spectrum (solid black line with black dots) at ${\\rm C}_2$ Swan (0,1) band head 5635.5~{\\AA}, plotted together with synthetic spectra with [C/Fe] values lowered by $-0.2$~dex (lower gray line) and $-0.25$~dex (upper gray line). } \\label{} \\end{figure}  The interpretation of the carbon abundance in NGC\\,6134 can be done by a comparison with carbon abundances determined for dwarf stars in the Galactic disk. Shi et al.\\ (2002) performed an abundance analysis of carbon for a sample of 90 F and G type main-sequence disk stars using C\\,{\\sc i} and [C\\,{\\sc i}] lines and found [C/Fe] to be about solar at the solar metallicity. The same result was found by Gustafsson et al.\\ (1999) who analysed a sample of 80 late F and early G type dwarfs using the forbidden [C\\,{\\sc i}] line. The ratios of [C/Fe] in our stars lie about 0.2~dex below the trend obtained for dwarf stars of the Galactic disk. Smiljanic et al.\\ (2009) analysed carbon abundances in two stars of the same cluster NGC\\,6134\\_30 and NGC\\,6134\\_202 and also found [C/Fe] abundance ratios to be lowered by about the same amount.  The interval 7980--8130~{\\AA}, with 11 CN lines selected, was analysed  in order to determine the nitrogen abundances. The mean nitrogen to iron abundance ratio in NGC\\,6134 is [N/Fe]=$0.25\\pm0.10$. Nitrogen to iron abundance ratios in the two clump  stars investigated in NGC\\,6134 by Smiljanic et al.\\ (2009)[N/Fe] ratios are enhanced  by 0.42 and 0.36~dex as well.  This shows that nitrogen is overabundant in these clump stars of NGC\\,6134, while [N/Fe]  values in the main-sequence stars are about solar at the solar metallicity  (c.f. Shi et al.\\ 2002). Reddy et al.\\ (2003) investigated nitrogen abundances in a sample  of 43 F--G dwarfs in the Galactic disk by means of weak N\\,{\\sc i} lines. At a value of  [Fe/H] of about $-0.2$~dex, which was well represented in their sample, [N/Fe] is about 0.2~dex.  There were few stars of solar metallicity investigated in this study. Nevertheless, the authors  make the extrapolation that at solar metallicity [N/Fe] values should be solar.  C/N ratios in our programme stars of NGC\\,6134 range from $0.98$ to $1.48$. For the stars NGC\\,6134\\_30 and NGC\\,6134\\_202, Smiljanic et al.\\ also found C/N to be $0.99$ and $1.41$, respectively.  \\input epsf \\begin{figure*} \\epsfxsize=\\hsize \\epsfbox[-20 -20 650 260]{Mik-fig6.eps} \\caption {Small region of NGC\\,6134\\_75 spectrum (solid black line with black dots) at ${\\rm C}_2$ Swan (1,0) band heads 4732.8~{\\AA} and 4735.3~{\\AA}, plotted together with synthetic spectra with [C/Fe] values lowered by $-0.2$~dex (lower gray line) and $-0.25$~dex (upper gray line). } \\label{} \\end{figure*}  The solar carbon and nitrogen abundances used in our work are log$A_{\\rm C} = 8.52$ and log$A_{\\rm N} = 7.92$ (Grevesse \\& Sauval 2000), so the solar ${\\rm C/N} = 3.98$. Smiljanic et al.\\ (2009) presented carbon and nitrogen values for ten open clusters in total and found an average C/N ratio of about $0.98$.   \\input epsf \\begin{figure} \\epsfxsize=\\hsize \\epsfbox[-20 -20 620 500]{Mik-fig7.eps} \\caption {Fit to the forbidden [O\\,{\\sc i}] line at 6300~{\\AA} in NGC\\,6134\\_39. The observed spectrum is shown as a solid line with black dots. Synthetic spectra with [O/Fe]$=-0.3$, $-0.25$, and $-0.2$ are shown as solid gray lines.} \\label{Oxygen} \\end{figure}  The analysis of oxygen abundance was performed using the most popular forbidden [O\\,{\\sc i}] line at 6300~{\\AA}. In the spectra of NGC\\, 6134 stars this line is not contaminated by telluric lines. In Fig.~7 we show an example of spectrum syntheses for [O\\,{\\sc i}] line in NGC\\,6134\\_39 star. This line is blended by Ni\\,{\\sc i}. The atomic data of Ni\\,{\\sc i} were taken from Johansson et al.\\ (2003).  The oxygen abundance results as a function of galactocentric distance are shown in Fig.~4. Similarly to other open clusters, [O/Fe] in NGC\\,6134 stars is lower than the solar value by about 0.2~dex. Oxygen abundances were not investigated in NGC\\,6134 by Smiljanic et al.\\ (2009).  In Fig.~8, [O/Fe] are plotted as a function of C/N ratios. The values of NGC\\,6134 are in agreement with other studies. At smaller C/N ratios [O/Fe] are lower.  \\input epsf \\begin{figure} \\epsfxsize=\\hsize \\epsfbox[-20 -20 620 500]{Mik-fig8.eps} \\caption { The oxygen abundance as a function of carbon and nitrogen abundance ratio. NGC\\,6134 stars are shown as filled squares, stars from open clusters M~67 and NGC\\,7789 investigated by Tautvai\\v{s}ien\\.e et al.\\ (2000, 2005) as circles, results by Smiljanic et al.\\ (2009) -- triangles. The solid line shows the values predicted by Schaller et al.\\ (1992).} \\label{Oxygen} \\end{figure}   \\subsection{Carbon isotope ratios}  The $^{12}{\\rm C}/^{13}{\\rm C}$ ratios were determined for all programme stars from the (2,0) $^{13}{\\rm C}^{14}{\\rm N}$ feature at 8004.728~{\\AA} with a laboratory wavelength adopted from Wyller (1966). In Fig.~9 we show a small portion of NGC\\,6134\\_157 spectrum together with three theoretical lines for carbon isotopic ratio determination. We find that the mean $^{12}{\\rm C}/^{13}{\\rm C}$ ratios are lowered to about  $9\\pm2.5$ in the clump stars investigated. Smiljanic et al.\\ (2009) found for this ratio the value of 12 and 13 for their two stars in NGC 6124. The $^{12}{\\rm C}/^{13}{\\rm C}$ ratio in the solar photosphere is equal to 89 (Coplen et al.\\ 2002).  The number of papers with carbon isotope ratios determined for red clump stars in open clusters is not numerous. In Fig.~10 we plot the mean carbon isotope ratios of clump stars in different open clusters as a function of turn-off mass and compare them with the theoretical models of the $1^{st}$ dredge-up by Boothroyd \\& Sackman (1999) and Charbonnel (1994).  For NGC\\,6134 we plot the mean $^{12}{\\rm C}/^{13}{\\rm C}$ ratio $10\\pm 3$ as averaged from clump stars investigated in our study and by Smilanic et al.\\ (2009). The mass of the programme stars $M = 2.34M_{\\odot}$ was obtained by Carretta et al.\\ (2004) reading the turn-off values on the Girardi et al.\\ (2000) isochrones for solar metallicity at the age of the cluster of 0.7~Gyr as determined by Bruntt et al.\\ (1999).  In Fig.~10, we also present results for clump stars in 5 other open clusters investigated by Smiljanic et al.\\ (2009). Data for M\\,67 and NGC\\,7789 clump stars come from Tautvai\\v{s}ien\\.{e} et al.\\ (2000 and 2005). From Gilroy (1989) we selected 4 clusters with well defined red clump stars (NGC\\,752, NGC\\,2360, M\\,67, and IC\\,4756). Luck (1994) derived carbon isotope ratios for 8 open clusters, however it is very difficult to identify stars of red clump in them.  From Fig.~10, it is seen that $^{12}{\\rm C}/^{13}{\\rm C}$ values in the clump stars are much below the predictions of $1^{st}$ dredge-up, and for most of the clusters they are even below the Cool Bottom Processing (CBP) model of extra-mixing (Boothroyd \\& Sackman 1999).  \\input epsf \\begin{figure} \\epsfxsize=\\hsize \\epsfbox[0 0 620 500]{Mik-fig9.eps} \\caption {Small region of NGC\\,6134\\_157 spectrum (solid black line with black dots) with $^{13}{\\rm C}^{14}{\\rm N}$ feature. Grey lines show synthetic spectra with $^{12}{\\rm C}$/$^{13}{\\rm C}$ ratios equal to 10 (lower line), 12 (middle line) and 14(upper line). } \\label{Oxygen} \\end{figure}    \\input epsf \\begin{figure} \\epsfxsize=\\hsize \\epsfbox[-20 -20 620 500] {Mik-fig10.eps} \\caption {The averaged carbon isotope ratios in clump stars of open clusters as a function of stellar  turn-off mass. $^{12}{\\rm C}$/$^{13}{\\rm C}$ in NGC\\,6134 (as averaged from this work and Smiljanic et al.)  is shown as a filled square, in M~67 and NGC\\,7789 as open squares (Tautvai\\v{s}ien\\.{e} et al.\\ 2000, 2005). Data for clusters investigated by Smiljanic et al.\\ (2009) are shown as triangles, and by Gilroy (1989) as plus  signs. The thick solid curve presents the theoretical model of $1^{st}$ dredge-up for  $Z=0.02$ by Boothroyd \\& Sackman (1999) and the thin solid line the one by Charbonnel (1994). The CBP model of  extra-mixing by Boothroyd \\& Sackman (1999) is plotted as a thick dashed line. } \\label{Oxygen} \\end{figure}   \\subsection{Sodium and Aluminum}  Sodium and aluminum are among the chemical elements for which observations of abundance anomalies are also present. The O-Na anticorrelation has been observed among the brightest red giants in Galactic globular clusters for a long time (see Kraft 1994; Da Costa 1998; Denissenkov \\& Herwig 2003 and references therein). An overabundance of Na could appear, due to the deep mixing from layers of the NeNa cycle of H burning. Extensive theoretical studies of deep mixing in stellar atmospheres have been made by Denissenkov \\& Weiss (1996), Denissenkov \\& Tout (2000), Denissenkov \\& Herwig (2003), Gratton et al.\\ (2004) and references therein. However, the explanation of abundance changes by deep mixing has been eliminated by the determination of an Na--O anticorrelation in less evolved stars down to the main sequence (Gratton et al.\\ 2001; Thevenin et al.\\ 2001; Ramirez \\& Cohen 2002; D'Orazi et al. 2010). For a recent discussion of the Na and O abundances in open clusters, see De Silva et al. (2009), who explicitly addressed the problem of the (not seen) Na--O anticorrelation.  Abundances of sodium were determined from the NLTE analysis of Na\\,{\\sc i} lines at 5682.64~{\\AA}, 6154.23~{\\AA} and 6160.75~(\\AA). The Na\\,{\\sc i} line at 5682.64~{\\AA} was not available for the analysis of UVES spectra. Abundances of aluminum were determined from Al\\,{\\sc i} lines at 7835.30~\\AA\\  and 7836.13~{\\AA}.  The stars in our sample do not show overabundance either of sodium, or of aluminum (Fig.~4).   \\subsection{The $\\alpha$-elements, iron group and heavier elements}  According to observations of main sequence stars in the Galactic disk, abundance ratios of $\\alpha$-process elements to iron at the solar metallicity are solar or slightly higher. In the study by Edvardsson et al.\\ (1993) [Mg/Fe], [Si/Fe] and [Ti/Fe] for almost all of the stars lie slightly above the solar ratio, [Ca/Fe] are solar. In the study of Reddy et al.\\ (2003), [Mg/Fe] and [Si/Fe] values are above solar, while [Ca/Fe] and [Ti/Fe] are exactly solar. The mean [$\\alpha$/Fe] ratios in a majority of open clusters investigated are slightly higher than in the Sun, and especially the overabundance of silicon is noticeable (c.f. Brown et al.\\ 1996; Bragaglia et al.\\ 2001; Friel et al.\\ 2003; but see also Bragaglia et al. 2008, Sestito et al. 2008). There has also been a claim that [$\\alpha$/Fe] is higher for outer disk clusters (e.g., Yong et al. 2005), but this has not been confirmed by other studies (Carraro et al. 2007, Sestito et al. 2008).  In NGC\\,6134, the mean cluster [$\\alpha/{\\rm Fe}] \\equiv {1\\over 4} ([{\\rm Mg}/{\\rm Fe}]+[{\\rm Si}/{\\rm Fe}]+[{\\rm Ca}/{\\rm Fe}] +[{\\rm Ti}/{\\rm Fe}]) = 0.02\\pm0.03$~(s.d.), which is very close to the solar value.  As it is seen from Table~3, the ratios of abundances of iron group, $s-$ ,and $r-$process elements to iron are close to solar.   \\subsection{Final remarks}  Convection, the only mechanism of internal mixing taken into account by standard stellar evolution models, is not able to account for carbon and nitrogen abundance alterations seen in clump stars of open clusters. The clump stars provide information on chemical composition changes, which have happened during their evolution along the giant branch and during the helium flash. Extra-mixing processes may become efficient on the red giant branch when stars reach the so-called RGB  bump, and may modify the surface abundances. It is known that alterations of $^{12}{\\rm C}/^{13}{\\rm C}$ and $^{12}{\\rm C}/^{14}{\\rm N}$ ratios depend on stellar evolutionary stage, mass and metallicity (see Charbonnel et al.\\ 1998, Gratton et al.\\ (2000), Chanam\\'{e} et al.\\ 2005, Smiljanic et al.\\ 2009 for more discussion).  In the open cluster M\\,67 (Tautvai\\v{s}ien\\.{e} et al.\\ 2000), we did not find evidence for extra-mixing in first-ascent giants (the mass of turn-off stars in this cluster is about 1.2~$M_{\\odot}$). In the M\\,67 giants investigated, the mean $^{12}{\\rm C}/^{13}{\\rm C}$ ratio is lowered to the value of $24 \\pm 4$ and the $^{12}{\\rm C}/^{14}{\\rm N}$ ratio to the value of $1.7 \\pm 0.2$, which is close to the corresponding predictions of the first dredge-up (Boothroyd \\& Sackmann 1999). Evidence of extra-mixing has been detected only in the clump stars observed, where the mean $^{12}{\\rm C}/^{13}{\\rm C}$ ratio is lowered to the value of $16 \\pm 4$, and the $^{12}{\\rm C}/^{14}{\\rm N}$ ratio to  $1.4 \\pm 0.2$. In  giants and clump stars investigated by Tautvai\\v{s}ien\\.{e} et al.\\ (2005) in NGC\\,7789, $^{12}{\\rm C}/^{13}{\\rm C}$ ratios are about the same, $9\\pm 1$, but $^{12}{\\rm C}/^{14}{\\rm N}$ ratio is $1.9\\pm 0.5$ in giants and $1.3\\pm 0.2$ in clump stars.  In NGC\\,6134 we investigated only clump stars. The mean $^{12}{\\rm C}/^{14}{\\rm N}$ ratio in these clump stars, $1.2\\pm 0.2$, and the mean $^{12}{\\rm C}/^{13}{\\rm C}$ ratio $9\\pm 2.5$, indicate a large extra-mixing, which is even larger than it is foreseen by the CBP model (Boothroyd \\& Sackmann 1999) for stars of a similar turn-off mass. $^{12}{\\rm C}/^{13}{\\rm C}$ ratios for the majority of clump stars in other open clusters also lie below the CBP trend (Gilroy 1989; Tautvai\\v{s}ien\\.{e} et al.\\ 2000, 2005; Smiljanic et al.\\ 2009). Further studies are required both on the theoretical and observational sides in order to improve our understanding of mixing processes in low and intermediate mass giants."
    },
    "1005/1005.4080_arXiv.txt": {
        "abstract": "We identify the L dwarf 2MASS~J20261584$-$2943124 as an unresolved spectral binary, based on low-resolution, near-infrared spectroscopy from IRTF/SpeX. The data reveal a peculiar absorption feature at 1.6~$\\micron$, previously noted in the spectra of other very low-mass spectral binaries, which likely arises from overlapping FeH and {\\meth} absorption bands in the blended light of an L dwarf/T dwarf pair. Spectral template matching analysis indicates component types of L0.5 and T6, with relative brightness $\\Delta{H}$ = 4.2$\\pm$0.6. Laser guide star adaptive optics imaging observations with Keck/NIRC2 fail to resolve the source, indicating a maximum separation at the observing epoch of 0$\\farcs$25, or a projected separation of 9 AU assuming a distance of 36$\\pm$5~pc. With an age that is likely to be relatively older ($\\gtrsim$5~Gyr) based on the system's large {\\vtan} and mass ratio arguments, the relative motion of the potentially ``massive'' (0.06-0.08~{\\msun}) components of 2MASS~J2026$-$2943 may be detectable through radial velocity variations, like its earlier-type counterpart 2MASS~J03202839$-$0446358 (M8+T5), providing dynamical mass measurements that span the hydrogen burning limit. ",
        "introduction": "Multiple systems are of fundamental importance in stellar astrophysics and are particularly crucial for studies of very low mass stars (VLM; M $<$ 0.1 M$_{\\odot}$) and brown dwarfs, the M, L, and T dwarfs (see \\citealt{2005ARA&A..43..195K} and references therein). The frequency and characteristics of VLM multiples provide empirical constraints for currently debated formation mechanisms (e.g., \\citealt{2007prpl.conf..427B, 2007prpl.conf..443L}), and are the dominant outlet for direct mass and radius measurements (e.g., \\citealt{2001ApJ...560..390L, 2004ApJ...615..958Z, 2006Natur.440..311S, 2008ApJ...689..436L, 2009ApJ...692..729D}). Resolved VLM multiples provide useful laboratories for studying atmospheric properties independent of age and composition effects (e.g., \\citealt{2005ApJ...634..616L, 2007ApJ...657.1064M, 2010ApJ...710.1142B}), and yield unique insights into evolutionary processes such as the poorly-understood L dwarf/T dwarf transition (e.g., \\citealt{2002AJ....124.1170D,2006ApJ...647.1393L, 2007ApJ...659..655B, 2008ApJ...685.1183L}).  Current VLM multiplicity studies are dominated by resolved imaging programs, which have uncovered $\\sim$100 systems to date\\footnote{For an updated list, see \\url{http://www.vlmbinaries.org}.}. These systems have constrained the frequency ($\\sim$15-35\\%), and separation and mass ratio distributions of VLM multiples. Yet resolved imaging studies are insensitive to very closely-separated ($\\rho <$ 1 AU) and/or distant multiples, as well as multiples observed in unfortunate geometries (e.g., \\citealt{1999Sci...283.1718M,2005ApJ...634..616L, 2006PASP..118..611G}). The closely-separated systems are particularly valuable, as they are the ones most likely to exhibit measurable orbital motion in a reasonable time frame, and also have the best chance of eclipsing.  Furthermore, there is growing evidence that a significant fraction - perhaps half - of VLM multiples are hidden in unresolved, tight systems, based on searches for radial velocity (RV) variables (e.g., \\citealt{2005MNRAS.362L..45M,2006AJ....132..663B,2008A&A...492..545J}) and overluminous sources in cluster color-magnitude diagrams (e.g., \\citealt{2003MNRAS.342.1241P,2007MNRAS.380..712L}).  Yet actual yields of very tight binaries have been small due to the inherent inefficiencies associated with these search programs: RV monitoring programs require large allocations of large telescope time, while searches for overluminous sources need clusters and long-term astrometric programs to confirm object membership.  Other signatures of unresolved multiplicity, such as eclipsing (e.g., \\citealt{2006Natur.440..311S}), astrometric wobble (e.g., \\citealt{2008ApJ...686..548D}) and microlensing (e.g., \\citealt{2009arXiv0911.2706B}) have similarly yielded few detections.   We have recently developed a technique to identify and characterize a subset of VLM ``spectral'' binaries, or blended-light pairs, using low-resolution, near-infrared spectroscopy \\citep{2007AJ....134.1330B, 2008ApJ...681..579B, 2010ApJ...710.1142B}. This technique exploits the distinct and complex structure of late-type M, L, and T dwarf near-infrared spectra as shaped by absorption bands of {\\wat}, {\\meth}, CO, and FeH, as well as condensate cloud opacity (e.g., \\citealt{2001AJ....121.1710R, 2003ApJ...596..561M, 2005ApJ...623.1115C}).  Unresolved pairs with different component spectral types exhibit specific spectral peculiarities which arise from blended features.  These peculiarities serve to both identify a source as a spectral binary and enable the decomposition of the spectrum into its component types.  This technique is not limited by the angular separation of a binary, so closely-separated systems---for which dynamic mass measurements and possibly radius measurements are feasible---can potentially be identified.  An illustrative case is 2MASS~J03202839$-$0446358 (hereafter 2MASS~J0320$-$0446; \\citealt{2003IAUS..211..197W}), whose near-infrared spectrum exhibits an unusual absorption feature at 1.6~$\\mu$m not seen in normal M or L dwarf spectra.  \\citet{2008ApJ...681..579B} were able to reproduce the spectrum of this source as the combination of an M8 primary plus T5 secondary, with the 1.6~$\\mu$m feature arising from overlapping FeH and {\\meth} absorption from the two components, respectively.  2MASS~J0320$-$0446 was independently identified as a spectroscopic binary with an 8-month orbit by \\citet{2008ApJ...678L.125B} from high-resolution spectroscopic monitoring.  The inferred separation and system mass function from the Blake et al.\\ study, and the component spectral types from the Burgasser et al.\\ study, have enabled a firm constraint on the minimum age and component masses of 2MASS~J0320$-$0446 ($>$2~Gyr; M$_1 >$ 0.080~{\\msun}, M$_2 >$ 0.053~{\\msun}; \\citealt{2009AJ....137.4621B}). To date, 20 late-type M/L plus T dwarf spectral binaries have been identified in this manner \\citep{2004ApJ...604L..61C, 2007AJ....134.1330B, 2008ApJ...681..579B, 2010ApJ...710.1142B, 2008ApJ...685.1183L, 2008arXiv0811.0556S}, three of which have thus far been resolved (\\citealt{2006ApJS..166..585B}; C.\\ Gelino et al. 2010, in preparation).  This article reports the discovery of a new L dwarf plus T dwarf spectral binary identified from low-resolution, near-infrared spectroscopy, {\\name} (hereafter {\\namesh}; \\citealt{2007AJ....133..439C}). The spectrum of this source shows the distinct signature of blended FeH and {\\meth} absorption, but is unresolved in high angular resolution imaging. In Section~2 we describe our spectroscopic observations of {\\namesh}, obtained with the NASA Infrared Telescope Facility (IRTF) SpeX spectrograph \\citep{2003PASP..115..362R}, and present our template fitting analysis that demonstrates the binary nature of this source and constrains its component properties. In Section~3 we describe high angular resolution images of {\\namesh} obtained with the Keck laser guide star adaptive optics (LGS AO) system and NIRC2 camera, and present limits on the detection of the putative companion. Section~4 provides additional discussion on the properties of the putative components, and presents arguments that the system may be relatively old ($\\gtrsim$5~Gyr). Results are summarized in Section~5. ",
        "conclusions": "We have identified {\\namesh} as a L0.5 plus T6 spectral binary, based on the presence of a peculiar 1.6~$\\micron$ feature in its near-infrared spectrum. We associate this feature with overlapping FeH and {\\meth} absorption from the primary and secondary, respectively. LGS AO imaging failed to resolve a companion to a minimum separation limit of 0$\\farcs$25, or 9~AU in projected separation, a fairly unrestrictive limit given the properties of VLM binaries at large. The inability to resolve this pair may be due in part to the large inferred magnitude difference of its components; its predicted $\\Delta{H}$ = 4.2$\\pm$0.6~mag is comparable to the largest relative magnitudes measured for any VLM pair (e.g., 2MASS~J1207334$-$393254AB, $\\Delta{H} \\approx 5.6$; \\citealt{2004A&A...425L..29C}; GJ~802AB, $\\Delta{H} \\approx 4.7$; \\citealt{2006ApJ...650L.131L}; SCR~J1845$-$6357, $\\Delta{H} \\approx 4.2$; \\citealt{2006ApJ...641L.141B}). Alternately, {\\namesh}, like 2MASS~J0320$-$0446, may be a very tightly bound system, raising the distinct possibility that direct mass measurements via astrometric and/or radial velocity orbital monitoring are viable. As few direct mass measurements are currently available for VLM stars and brown dwarfs, additional follow-up observations---e.g., LGS AO aperture masking observations \\citep{2006SPIE.6272E.103T} or high resolution spectroscopic monitoring---to resolve this system and/or measure the properties of its components is warranted. {\\namesh} joins a growing list of unresolved, VLM binary candidates whose inferred components straddle the two lowest-luminosity classes of stars and brown dwarfs."
    },
    "1005/1005.3496_arXiv.txt": {
        "abstract": "{% The disciplines of asteroseismology and extrasolar planet science overlap methodically in the branch of high-precision photometric time series observations. Light curves are, amongst others, useful to measure intrinsic stellar variability due to oscillations, as well as to discover and characterize those extrasolar planets that transit in front of their host stars, periodically causing shallow dips in the observed brightness. Both fields ultimately derive fundamental parameters of stellar and planetary objects, allowing to study for example the physics of various classes of pulsating stars, or the variety of planetary systems, in the overall context of stellar and planetary system formation and evolution. Both methods typically also require extensive spectroscopic follow-up to fully explore the dynamic characteristics of the processes under investigation. In particularly interesting cases, a combination of observed pulsations and signatures of a planet allows to characterize a system's components to a very high degree of completeness by combining complementary information. The planning of the relevant space missions has consequently converged with respect to science cases, where at the outset there was primarily a coincidence in instrumentation and techniques. Whether space- or ground-based, a specific type of stellar pulsations can themselves be used in an innovative way to search for extrasolar planets. Results from this additional method at the interface of stellar pulsation studies and exoplanet hunts in a beyond-mainstream area are presented.% } ",
        "introduction": "\\label{sec:introduction} The (transiting) extrasolar planet fields and the asteroseismology field see a convergence of instrumentation that culminates in the insight that beyond this purely technical level, a much more fundamental connection exists in the shared desire for the most exhaustive characterization of stellar and planetary systems at all possible with the available diagnostics. I will first sum up the relevant current context with an emphasis on planets around evolved stars, and then specifically address the topic of oscillation timing as a means to detect planets. \\par In terms of successes to detect planets around evolved stars, the subdwarf B stars as host stars stand out as a group. This class of evolved objects will be introduced, and the difficulties in explaining their sheer existence mentioned. Incidentally, the asteroseismology of subdwarf B stars is also a very active field. The \\mbox{EXOTIME} planet searching program will be presented, which takes advantage of the long-term behaviour of these pulsations. Current ideas on subdwarf B evolution, and the potentially crucial role of planets, will be discussed, frequently resorting to the \\object{V391 Pegasi} system. \\begin{figure*} \\begin{center} \\includegraphics[width=\\textwidth,angle=0]{macp.eps} \\end{center} \\caption{The ''Perryman tree'': Detection methods for extra-solar planets with detectable masses on a (logarithmic) mass scale. Adopted from Figure~1 in \\citet{2000RPPh...63.1209P}, updated to include recent detections up to February 2010 (courtesy of M.\\ A.\\ C.\\ Perryman). } \\label{fig:perrymantree} \\end{figure*} ",
        "conclusions": "\\label{sec:summary} This article has attempted to highlight the added value of incorporating stellar pulsations in the comprehensive investigation of various stellar and planetary systems. \\par The potential of links between asteroseismology and\\linebreak exoplanet science was first illustrated for solar-like oscillators. Accurate asteroseismic radii for (solar-like) host stars of transiting exoplanets translate into good radius determinations for the transit planets. \\par In the context of connections with coherent pulsations, as found in evolved stars, the focus was directed towards the investigation of stellar and planetary systems at late stages of stellar evolution. The asteroseismology-exoplanet-con\\-nection was explored in particular for post-RGB stars. The exotic class of subdwarf B stars, of which many show pulsations, was introduced as a group with several new planet discoveries. In discussing the fate of planetary systems based on known planets around extreme horizontal branch stars (e.g. \\object{V391 Pegasi}), the question arises whether these are first or second generation planets. Data from programs such as \\mbox{EXOTIME}, which make use of the timing method to find more such systems, can fill a gap with respect to the sensitivity to planets in wide orbits and applicability to evolved stars. \\par With respect to the puzzle of subdwarf B formation, planets were proposed to (partly) be at the origin of sub\\-dwarf B stars, and also to play a role as the second parameter for the horizontal branch morphology in globular clusters and elliptical galaxies. \\par The aspects this article summarizes are all from the pre-\\textit{Kepler} era. It will be fascinating to see how our understanding in a great many areas evolves with the new data."
    },
    "1005/1005.5708_arXiv.txt": {
        "abstract": "The small-scale magnetic helicity produced as a by-product of the large-scale dynamo is believed to play a major role in dynamo saturation. In a mean-field model the generation of small-scale magnetic helicity can be modelled by using the dynamical quenching formalism. Catastrophic quenching refers to a decrease of the saturation field strength with increasing Reynolds number. It has been suggested that catastrophic quenching only affects the region of non-zero helical turbulence (i.e.\\ where the kinematic $\\alpha$ operates) and that it is possible to alleviate catastrophic quenching by separating the region of strong shear from the $\\alpha$ layer. We perform a systematic study of a simple axisymmetric two-layer $\\alpha\\Omega$ dynamo in a spherical shell for Reynolds numbers in the range $1 \\leq \\Rm \\le 10^5$. In the framework of dynamical quenching we show that this may not be the case, suggesting that magnetic helicity fluxes would be necessary. ",
        "introduction": "It is widely believed that the solar magnetic cycle is driven by an $\\alpha\\Omega$ dynamo. The region of strong radial shear called the tachocline \\citep{SZ92} at the bottom of the solar convection zone (SCZ) is believed to be the place where strong toroidal field is formed due to stretching of the weaker but diffuse poloidal field. It has been inferred from helioseismology that the tachocline is confined within a thin layer in the overshoot region which lies below the SCZ. This led \\cite{Par93} to propose the idea of an interface dynamo, where the shear is confined to a region with a greatly reduced turbulent diffusion, which also is the region of production of strong toroidal field. The helical turbulence generated due to convection and rotation in the layer above provides the turbulent $\\alpha$ and a large turbulent diffusivity $\\eta_{\\rm t}$. The dynamo cycle being thus completed, the interface dynamo operates as a surface wave propagating along the boundary between strong shear and convection, which is also a region with a strong gradient in the turbulent diffusivity.   In order to mimic the process of dynamo saturation, traditionally authors have used an algebraic quenching function like $\\alpha_0/(1+\\overline{\\bm{B}}^2/B_{\\rm eq}^2)$ or even $\\alpha_0/(1+R_{\\rm m} \\overline{\\bm{B}}^2/B_{\\rm eq}^2)$, where $\\alpha_0$ is the unquenched value of $\\alpha$, $\\Rm$ is the magnetic Reynolds number, $\\overline{\\bm{B}}$ is the mean magnetic field and $B_{\\rm eq}$ is the equipartition magnetic field. However, it appears that conservation of magnetic helicity plays an important role in the process of saturation. At large $\\Rm$, the magnetic helicity ($\\int\\bm A\\cdot\\bm B\\, dV$) is fairly well conserved by the dynamo producing equal amounts of helicity in small and large scales, respectively. If the small-scale magnetic helicity is not able to escape out of the system, then the turbulent  $\\alpha$ effect is markedly reduced \\citep{PFL}. This leads to the magnetic energy of the dynamo to be quenched such that the saturation value varies as $\\Rm^{-1}$. Since astrophysical objects have large $\\Rm$ ($\\Rm\\sim10^9$ for the Sun), this strong dependence of the saturation field strength $B_{\\rm sat}$ on $\\Rm$ is referred to as {\\em catastrophic quenching}.   Even though the helicity constraint in direct numerical simulations (DNS) of dynamos with strong shear has been clearly identified, the results can be matched with mean-field models having a weaker algebraic quenching of $\\alpha$ and turbulent diffusivity than $\\alpha^2$ dynamos \\citep{BBS01}. However, the empirically determined coefficients would depend on circumstances and are therefore not universal. Such a model would not obey magnetic helicity evolution and is therefore untenable on theoretical grounds. The interface dynamo model has been invoked several times as a way of getting around the persistent problem of  catastrophic quenching \\citep[see][for a review]{Cha05}. The belief is that the quenching function remains close to unity in the region of finite $\\alpha$ since the toroidal field is expected to be weak there. However, to our knowledge, this has never been verified in a consistent manner for a range of magnetic Reynolds numbers. In this paper we perform a series of calculations with mean-field $\\alpha\\Omega$ models, in spherical geometry, considering both algebraic and dynamical quenching formulations, for magnetic Reynolds numbers in the range $1\\le \\Rm \\le 2\\times10^5$. An important feature of these models is that the region of strong narrow shear is separated from the region of helical turbulence as proposed in the Parker's interface model.  In \\S2 we discuss the features of the $\\alpha\\Omega$ model used, and the formulation of  dynamical $\\alpha$ quenching. The results are highlighted in \\S3 and  conclusions are drawn in \\S4.  A part of the calculations presented in this paper will be discussed in a more detailed paper \\citep{CGB10}. In this paper we focus specifically on the reality of the catastrophic quenching in dynamos with $\\alpha$ and $\\Omega$ effects operating in two widely separated layers. ",
        "conclusions": "The calculations done in this paper indicate that it is not possible to escape catastrophic quenching due to accumulation of small-scale helicity in the domain by merely separating the regions of shear and $\\alpha$ effect. The saturation value of magnetic energy decreases as $\\sim R_{\\rm m}^{-1}$ for both dynamical quenching and the `strong' ($\\sim R_{\\rm m}$-dependent) form of algebraic quenching for the simple two-layer model. However, additionally we observe parity fluctuations for cases with dynamical quenching. It does not seem to us that there exists any chaotic behaviour in the time series of magnetic energy since the dynamo period and the saturation energy remains fairly constant. It may be possible that solar wind, coronal mass ejections, and Vishniac and Cho fluxes help in throwing out the small scale helicity from the Sun and thus alleviate catastrophic quenching. In this study we have not found any difference between the nature of the saturation curves for an $\\alpha\\Omega$ dynamo and an $\\alpha^2$ dynamo using the form of dynamical quenching given by Eq.~(\\ref{eq:alphaeq}). However, it is clear that the algebraic quenching formula must fail if we were to allow for magnetic helicity fluxes that would, under suitable circumstances, alleviate catastrophic quenching.  We have been cautious about using dynamical quenching equation for dynamo numbers not very large compared to the critical dynamo numbers. We would expect that the magnetic field should affect all the turbulence coefficients including both $\\alpha$ and $\\eta$. However for this analysis we have not included any quenching of $\\eta_{\\rm t}$. This may be justified since such an effect could be mimicked by our simple two-layer model with a lower $\\eta_{\\rm t}$ in the region of production of strong toroidal fields and a higher $\\eta_{\\rm t}$ in the region of weaker poloidal fields. The effect of dynamical quenching on more realistic solar dynamo models having meridional circulation, Babcock-Leighton $\\alpha$ effect, and diffusive helicity fluxes have also been studied \\citep{CGB10}.  Unfortunately the direct numerical simulations have not yet reached the modest Reynolds numbers used in this paper ($\\sim 10^4$) which are still much lower than the astrophysical dynamos. If $B_{\\rm sat}^2$ really has an inverse dependence on $\\Rm$, then the solar dynamo should not be operating like it does! We may conclude that either we must have helicity fluxes out of the system or cross-equatorial diffusive fluxes inside the domain or that the $\\alpha_{\\rm M}$ equation to be used beside the mean-field induction equations must be suitably modified for $\\alpha\\Omega$ dynamos. To verify if the equation for dynamical quenching works in the same way as in $\\alpha^2$ dynamos, we need to perform systematic comparisons between DNS with shear and convection and mean-field modelling for $\\alpha\\Omega$ dynamos."
    },
    "1005/1005.2128_arXiv.txt": {
        "abstract": "The Chandra Orion Ultradeep Project observed hundreds of young, low-mass stars undergoing highly energetic {\\sc x}-ray flare events. The 32 most powerful cases have been modeled by \\citet{Favata:2005} with the result that the magnetic structures responsible for these flares can be many stellar radii in extent. In this paper, we model the observed spectral energy distributions of these 32 stars in order to determine, in detail for each star, whether there is circumstellar disk material situated in sufficient proximity to the stellar surface for interaction with the large magnetic loops inferred from the observed X-ray flares. Our spectral energy distributions span the wavelength range 0.3--8~$\\mu$m (plus 24~$\\mu$m for some stars), allowing us to constrain the presence of dusty circumstellar material out to $\\gtrsim 10$~{\\sc au} from the stellar surface in most cases. For \\goodguys\\ of the 32 stars in our sample the available data are sufficient to constrain the location of the inner edge of the dusty disks.  \\catwcone\\ of these (\\catpone) have spectral energy distributions consistent with inner disks within reach of the observed magnetic loops. Another \\catwtwo\\ stars may have gas disks interior to the dust disk and extending within reach of the magnetic loops, but we cannot confirm this with the available data. The remaining \\catthree\\ stars (\\catpthree) appear to have no significant disk material within reach of the large flaring loops. Thus, up to $\\sim 40$\\% of the sample stars exhibit energetic {\\sc x}-ray flares that possibly arise from a magnetic star-disk interaction, and the remainder are evidently associated with extremely large, free-standing magnetic loops anchored only to the stellar surface. ",
        "introduction": "}  The recent large {\\sc x}-ray surveys of the Orion and Taurus star-forming regions performed by {\\it Chandra} and {\\it XMM} \\citep[i.e.,\\ {\\sc coup}, {\\sc xest}:][]{COUP:2005,XEST:2007} provide an unparalleled opportunity to study the magnetic activity of young, low-mass stars. These deep observations spanning long temporal baselines (e.g.,\\ the {\\sc coup} {\\sc x}-ray light curves span 13 days with near-continuous time coverage) reveal that low-mass pre-main-sequence ({\\sc pms}) stars possess {\\sc x}-ray luminosities 3--4 magnitudes greater than that of the present-day Sun and exhibit extremely energetic flaring events with high frequency.  Detailed analyses of these flares reveal that they are similar to solar flares, but are orders of magnitude more energetic and larger in physical size. In particular, \\citet{Favata:2005} subjected the 32 most energetic flares observed by {\\sc coup} to analysis via a standard uniform cooling loop model \\citep{Reale:1997,Sylwester:1993,Priest:2002,Favata:2003}, with which they derived the properties of the magnetic coronal loops that participate in the flare events. They found that these magnetic loops were extremely large---extending tens of stellar radii in some cases---much larger than ever observed on older stars. Such large-scale flares could have important ramifications for a number of issues, such as the shedding of stellar angular momentum and mass, powering of outflows, and ionization/dissipation of circumstellar disks.  Magnetic loops with sizes on the order of $\\sim10$ stellar radii have long been postulated as part of magnetospheric accretion scenarios. In this paradigm, a large-scale stellar magnetic field threads the inner edge of a circumstellar disk, channeling accretion from disk to star \\citep[e.g.,][]{Camenzind:1990,Konigl:1991,Shu:1994,Hartmann:1994,Hayashi:1996}. Indeed, \\citet{Favata:2005} speculated that the large magnetic loops observed in the {\\sc coup} sample may be facilitating this type of magnetic star-disk interaction, in part because they argued that such large loops would likely be unable to remain stable if anchored only to the stellar surface.  However, as most of the 32 {\\sc coup} sources studied by \\citet{Favata:2005} lacked at that time sufficient photometric data with which to characterize the optical-infrared ({\\sc ir}) spectral energy distributions, they could not confirm the presence of inner disks to which the observed magnetic loops might link.  Thus, there is still an outstanding question as to whether stellar coronal activity in these stars alone can drive such energetic flare events, or whether the energy (or at least the trigger) derives from a star-disk interaction.  While the latter requires magnetic loops large enough to reach the inner edge of the disk, several theoretical studies \\citep[e.g.,][and references therein]{Uzdensky:2002,Matt:2005a} have shown that the presence of a disk truncates the stellar magnetosphere so that closed magnetic loops extend not much further than the inner edge of the disk.  Thus, if these flares are powered by the star-disk interaction, one would expect the size of the flaring loops to approximately coincide with the location of the disk inner edge.  On the other hand, if the energetic flares are purely a stellar phenomenon, the largest loop sizes may only be exhibited by stars that lack disk material close to the star.  To address this question, it will therefore be useful to be able to determine the proximity of disk material to the furthest extent of the flaring magnetic loops exhibited by the sample analyzed in this work.  Near- and mid-{\\sc ir} colors can be used as a crude tracer of close-in circumstellar material. \\citet{Getman:2008b} have used {\\it Spitzer} {\\sc ir} colors to distinguish Class~{\\sc ii} and Class~{\\sc iii} objects (i.e., stars with dusty disks and naked T~Tauri stars, respectively) among 161 flaring {\\sc coup} stars. Interestingly, they found evidence that whereas the largest flaring loops tended to be associated with Class III objects, the Class~{\\sc ii} sources in their sample were more likely to possess relatively small magnetic loops. The {\\sc ir} colors alone do not provide a quantitative measure of the location of the inner disk edge, but \\citet{Getman:2008b} suggest that the magnetic loops may be confined by the disk to be within the disk co-rotation radius (the radius at which disk material, if present, orbits the star with angular velocity equal to the star's angular velocity). This suggestion is important, as the disk co-rotation radius is the point specifically at which some magnetospheric accretion theories predict magnetic star-disk interaction to occur \\citep[e.g.,][]{Ostriker:1995}. It is desirable, therefore, to establish the relationship between circumstellar disks and the very large magnetic loops observed by {\\sc coup} more quantitatively than {\\sc ir} colors alone permit.  In this paper, we present detailed spectral energy distributions ({\\sc sed}s) for each of the 32 most powerful {\\sc x}-ray flaring {\\sc coup} sources at wavelengths 0.34--8~$\\mu$m, plus upper limits at 24~$\\mu$m (\\S\\ref{data}). In \\S\\ref{excessive} we present near-infrared color excesses for the sample as a basic tracer of close-in circumstellar disks. Next we compare in detail the full observed {\\sc sed}s against synthetic {\\sc sed}s of low-mass {\\sc pms} stars with disks (\\S\\S\\ref{models}--\\ref{example}) in order to (a) ascertain whether dusty disks are present around these stars, and (b) if so, determine quantitatively whether the inner edges of those disks are sufficiently close to the stellar surface to interact with the large flaring loops observed by \\citet{Favata:2005}. The results (\\S \\ref{results}) indicate that approximately more than half of the sample stars lack significant disk material within reach of their flaring magnetic loops; evidently the extremely large flaring loops observed by \\citet{Favata:2005} are in most cases free-standing structures anchored only to the stellar surface. In \\S\\ref{discussion} we discuss some implications of this finding. ",
        "conclusions": "}  Of the \\goodguys\\ stars in our 32-star sample for which we have enough optical--infrared data to constrain the location of the disk inner edge (i.e., excluding stars in category 4; \\S\\ref{criteria} and Table~\\ref{summary}), for about \\catpthree\\ we are able to rule out close-in, massive disks within reach of the observed large flaring loops (category 3). These energetic flares discovered by \\citet{Favata:2005} are evidently intrinsically stellar phenomena. This gives added justification {\\it a posteriori} for the application of the solar-flare cooling loop model to these stars (\\S\\ref{data2p1}), and suggests that it may be possible by further extension of the solar analogy to infer other flare-related properties for these flares, such as coronal heating rates and coronal mass ejections.  The latter in particular may be important for furthering our understanding of mass and angular momentum loss in these low-mass {\\sc pms} stars.  Our sample also includes \\catwone\\ cases for which the {\\sc sed} clearly indicates a dusty disk that extends close enough to the star to permit interaction with the flaring loop (category 1; Table~\\ref{summary}).  In \\catwtwo\\ additional cases the dust disk appears to be truncated beyond the reach of the flaring loop, but at or close to the predicted dust destruction radius (category 2).  In these cases, the dust disk may in fact be truncated by dust sublimation, a process which does not remove gas.  Thus, in category 2 objects, it seems likely that a gas disk (undetected in the broadband flux measurements used in our {\\sc sed} models) extends closer to the star and may be within reach of the observed flaring loops.  One of the \\catwtwo\\ category 2 objects has a \\ion{Ca}{2} measurement from the literature ({\\sc coup} 141), and two have $\\Delta(U-V)$ measurements ({\\sc coup} 141 and {\\sc coup} 1568).  {\\sc coup} 141 is, interestingly, the most strongly accreting object in the sample as probed by its \\ion{Ca}{2} equivalent width (see Table~\\ref{stelparms}) and it also has a very negative $\\Delta(U-V)$; the combination of these indicators is strong evidence for ongoing accretion. For {\\sc coup} 141 the \\dik\\ near-{\\sc ir} excesses also indicates a disk very close to the star (see Sec.~\\ref{excessive} and Table~\\ref{foursquare}). These examples further strengthen the interpretation of the category 2 stars as likely having gas accretion disks within reach of the stellar magnetosphere.  Several studies clearly find a high frequency of close-in, massive dusty disks in the {\\sc onc} population as a whole. For example, \\citet{LAH:1998} find a disk fraction in the {\\sc onc} of $\\sim$70\\% on the basis of excess emission at 2.2~$\\mu$m. Thus, the $\\sim$\\catpone\\ (category 1) or $\\sim$38\\% (categories 1 and 2) frequency of close-in disks in our 32-star sample, representing the $\\sim$1\\% of {\\sc onc} stars with the most powerful X-ray flares observed by COUP, is evidently not representative of the disk characteristics of the {\\sc onc} as a whole.  The \\catwone\\ cases in which we have observed stars with close-in, dusty disks which intersect the magnetic loops (category 1) are interesting candidates for further study.  Specifically, it would be informative to determine if the flares in these cases are in some way different than the category 3 cases.  Three of these objects have \\duv\\ measurements, all of them $\\le -0.3$ (Table~\\ref{stelparms}), strongly indicating active accretion.  Temporally linking accretion as seen in optical variability to the X-ray flare events \\citep[e.g.,][]{Stassun:2006} could solidify whether a magnetic star-disk interaction has taken place. Geometric information would be necessary to determine where the magnetic loop is on the stellar surface---for example, a given flaring loop could extend in a direction perpendicular to the disk and thus not interact, even if the dust (or gas) disk is within the appropriate distance from the stellar surface.  The question of how the large magnetic structures are stabilized (prior to the flaring event) was posed in the discovery publication of these objects \\citep{Favata:2005}.  It was proposed that the loops may be anchored to corotating disk material and thereby not subject to shear which could disrupt the loops. In this work we have found such disk-anchoring to indeed be a possibility for 10 objects (categories 1 and 2), but further analysis is needed to determine how massive and ionized a disk must be to enable interaction.  For \\catwthree\\ objects (category 3), disk-supported loops are unlikely, as these stars lack massive disks within reach of the loops. We speculate that as long as the confining magnetic field at the upper end of the loop is sufficiently strong and the confined material corotating, stability is feasible even without a disk.   For example, \\citet{Cranmer:2009} describes a loop geometry in which the pressure of the confined gas decreases with increasing loop length, implying that in fact the largest coronal loops may be most stable against rupture. Additionally, our findings may imply that the largest flaring loops cannot readily form in the presence of a disk, given that they appear in our sample to preferentially occur on stars lacking close-in disk material.  Alternate scenarios to explain the long {\\sc x}-ray decay timescales include observations of corotating, embedded structures in coronae \\citep[e.g.,][]{Cameron:1989a,Cameron:1989b}, and prominences and stellar winds \\citep{Skelly:2008,Massi:2008}.  The observed coronal structures from those studies lend additional support to the idea of these large magnetic structures remaining stable in hot coronae or within the stellar wind over multiple rotation periods as the {\\sc x}-ray flare decays, even if no disk is present to anchor the magnetic loop \\citep{Jardine:2005}.  In summary, the 32 most powerful flares observed by the {\\sc coup} survey were found to have magnetic structures multiple stellar radii in arc length confining the {\\sc x}-ray emitting, heated plasma \\citep{Favata:2005}. With the goal of understanding the nature of these large {\\sc x}-ray emitting flare structures, we have modeled the optical--infrared {\\sc sed}s of these objects, finding \\catpthree\\ to be lacking close-in circumstellar disks to which these loops could anchor.  It is evident that in at least these cases the large-scale flares are phenomena of purely stellar origin, neither triggered nor stabilized by star-disk interactions."
    },
    "1005/1005.4823_arXiv.txt": {
        "abstract": "We present the integrated \\Ha\\ emission line profile for 157 \\HII\\ regions in the central $3.4'\\times 3.4'$ of the galaxy M\\,83 (NGC\\,5236). Using the Fabry-Perot interferometer \\ghafas, on the 4.2 m William Herschel on La Palma, we show the importance of a good characterization of the instrumental response function for the study of line profile shapes.  The luminosity-velocity dispersion relation is also studied, and in the $\\log(L)-\\log(\\sigma)$ plane we do not find a linear relation, but an upper envelope with equation $\\log(L_{{\\rm H}\\alpha})=0.9 \\cdot \\log(\\sigma)+38.1$. For the adopted distance of 4.5 Mpc, the upper envelope appears at the luminosity $L=10^{ 38.5}$ \\ergs, in full agreement with previous studies of other galaxies, reinforcing the idea of using \\HII\\ regions as standard candles. ",
        "introduction": "\\label{sec:intro}  \\HII\\ regions are among the brightest non-ephemeral objects in galaxies, with prominent emission lines characterizing their spectra. By studying the integrated line profiles of such emission lines we can, to some extent, infer the physical properties inside the regions, the virial mass of the system, temperatures, densities, and their ages. One important parameter for deriving dynamical properties is the velocity dispersion of the emission line $\\sigma={\\rm FWHM}/2\\sqrt{2\\ln 2}$, where FWHM is the full width at half maximum of the line. By studying the integrated \\Ha\\ emission-line profiles of giant extra-galactic \\HII\\ regions, \\cite{SmithandWeedman70} found velocity dispersions, $\\sigma \\sim 25$ \\kms. Since then, great efforts have been made to explain the line profiles of \\HII\\ regions. \\cite{Melnick77} and \\cite{TerlevichandMelnick81} analyzed 12 regions and claimed to find a relation between sizes, luminosities $L$, and integrated emission line velocity dispersion $\\sigma$ \\citep[for further discussion in this topic see][]{{Melnicketal88},{TellesandTerlevich93}}. One fundamental relation reported by these authors is $L \\propto \\sigma^4$, which is equivalent to the Faber-Jackson relation for elliptical galaxies \\citep{FaberandJackson76}. Given the potential importance of finding a valid direct relation between $L$ and $\\sigma$ to understand the internal dynamics of \\HII\\ regions but that observations have not seemed to yield a unique relation, it has been the subject of some controversy.  After an initial claim by \\cite{TerlevichandMelnick81} that they had observed a relation $L \\propto \\sigma^4$ implying that \\HII\\ regions are in virial equilibrium,  \\cite{Royetal86} and \\cite{Hippelein86} found a similar relation, but with exponents $\\sim 3$ and $\\sim 6$, respectively, while \\cite{GallagherandHunter83} found no clear proportionality between $L$ and $\\sigma$. On the other hand, \\cite{Melnicketal87} adduced problems in the selection of the sample made by Gallagher \\& Hunter, since they used regions with low velocity dispersion, while \\cite{TerlevichandMelnick81} used giant \\HII\\ regions, with high velocity dispersion ($\\sigma >13$ \\kms ) \\black. \\cite{Melnicketal87} reported a new value of $\\sim 5$ for the exponent in the $L-\\sigma$ relation. A common characteristic of all these studies is that they use at most few tens of the brightest regions, most of them in the high velocity dispersion range \\black, taken from different galaxies, i.e., introducing extra uncertainties from the assumed distances to different systems. This produces additional scatter as extinction and distances are different for each galaxy.  The obvious next step was then to study all the regions available from the same galaxy where the distance would have no effect on the scatter and the extragalactic and Galactic extinction would be similar. \\cite{Arsenaultetal90} studied 127 regions in the galaxy NGC\\,4321 and found no clear $L-\\sigma$ relation, except for the subsample of the highest surface brightness regions, for which they report an exponent of $\\sim 2.6$. Probably the most remarkable aspect of the study by \\cite{Arsenaultetal90} is that they define an upper envelope for the $\\log(L)-\\log(\\sigma)$ diagram, an upper limit for the luminosity at a certain velocity width. This envelope was also found by \\cite{Rozasetal98} and \\cite{Relanoetal05} with slopes of $2.6$ and $2.0$ respectively, and again by \\cite{Rozasetal06a} who found 2.9 for a sample of supergiant \\HII\\ regions. All these papers relate the upper envelope to where the regions in virial equilibrium are located, while broader line profiles have contribution from non-equilibrium processes   The line profiles of spatially unresolved \\HII\\ regions have been traditionally fitted by Gaussians  \\citep[][ and many others]{{SmithandWeedman70},{TerlevichandMelnick81}, {Arsenaultetal90}}. \\cite{Rozasetal06b} and \\cite{Westmoquetteetal07} distinguish a broad low intensity component added to the main spectral component.   This has been more a question of mathematical simplicity than of the physical processes which influence the line shape. Studies of spatially resolved \\HII\\ regions \\citep[e.g.][]{{MeaburnandWalsh81},{Meaburn84}, {Clayton86}, {ChuandKennicutt94}} showed clearly that not only do they receive multiple inputs of energy, but that these inputs vary from place to place (and time to time) within a region, so that the resulting integrated emission line profile is a luminosity weighted mean of all these inputs. We can pick out two principal contributors to the final line shape: the turbulent motion within the ionized gas (particularly where it interacts with the internal and surrounding neutral gas), and the emission from dominant expanding shells \\citep[][]{{Dyson79},{Meaburn80},{GallagherandHunter83}}. The contribution from the turbulent motions is gaussian, as it is produced by random motion, while the expanding shells, if present, are identified as equidistant in velocity from the main peak. These symmetric components produce non-gaussian wings or, if the expansion velocity and the luminosity are high enough, secondary peaks. In addition to these two contributions, there may be more local expansion due to the winds of multiple individual stars and supernova shells \\citep{Redmanetal03}. Also the hyperfine structure of Hydrogen split \\Ha\\ in seven components $\\sim$0.14 \\AA\\  apart, which should be considered.   With such a complex ensemble of contributions it is importantthe to take into account the shape of the spectral instrument response to accurately correct its effect on observations. Although, again for the sake of mathematical simplicity, it has traditionally been assumed Gaussian, the instrumental response is better described by a Lorentzian \\citep[e.g.][]{Bland-Hawthorn95}. Thus, instead of fitting a Gaussian (assuming the turbulent motion as the dominant feature in the line profile), \\cite{MoiseevandEgorov08} concluded that Voigt profiles are a more realistic characterization of the integrated emission line profiles of \\HII\\ regions, since it meant assuming the instrumental response is Lorentzian, not Gaussian (more about this will be discussed in sect.~\\ref{sec:fitting}).   During the present study we take particular care of the instrumental response and fit Gaussians convolved with it to the integrated \\Ha\\ line profiles, comparing it with the direct fit to a Gaussian. In sect.~\\ref{sec:observations} we describe the observations, sect.~\\ref{sec:reduction} explains the data reduction, focusing on the improved steps carried out for the present study. The line profiles are analyzed in sect.~\\ref{sec:analysis}, with special emphasis in the $L-\\sigma$ relation. Results and discussion are in sect.~\\ref{sec:results} and conclusions presented in \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} We have presented an improved method for analysing emission line profiles from Fabry-Perot interferometers, and have applied it to \\ghafas\\ Fabry-Perot data for the central 2.3 kpc radius of the nearby Sc galaxy M\\,83. We have considered the true instrumental profile without any parametrisation, and derived observed emission line profile parameters by a convolution of a Gaussian with the instrumental response.  From a narrow-band \\Ha\\ image, a catalog of \\HII\\ regions was extracted and we were able to determine the flux, position, and area for 1029 star forming \\HII\\ regions. Improving some of the Fabry-Perot data reduction procedures (sect.~\\ref{sec:reduction}) we have measured accurate integrated \\Ha\\ emission line profiles for 157 \\HII\\ regions in the \\ghafas\\ field.  Convolving with the IRC before fitting is the more realistic approach and clearly gives better fitting results in terms of smaller $\\chi^2$ for 76\\% of the regions. When fitting two Gaussians, the errors in the determination of the line parameters increase to more than 40\\%, due to low number of degrees of freedom, and for this reason we ruled out successful multiple-Gaussian fitting. Following the commonly used procedure, when fitting Gaussians in a direct way, the IRC was assumed a Gaussian and quadratically subtracted. The direct fitting method leads to an overestimation of $\\sigma$ by a constant value of $\\sim 8$ \\kms\\ . We used the results from our improved line fitting approach to study the $L-\\sigma$ relation and found no significant linear correlation between these two variables, but an upper envelope defined by $\\log(L_{{\\rm H}\\alpha})=1.3 \\cdot \\log(\\sigma)+37.6$. Our data is fully consistent with the previous estimations of \\cite{Arsenaultetal90}, \\cite{Rozasetal98} and \\cite{Relanoetal05}, when applying direct Gaussian fitting.   The upper envelope emerges in the $\\log(L)-\\log(\\sigma)$ diagram at luminosities higher than $10^{38.5}erg\\, s^{-1}$.  \\cite{Beckmanetal00} suggested that this luminosity might correspond to the transition from an ionization bounded regime to a density bounded regime (i.e. from optically thick to optically thin). We now understand, in part from work by \\cite{Giammancoetal04} that the simple Str\\\"omgren  sphere scenario on which this concept is based is too simple, and that the strongly inhomogeneous structure of \\HII\\ regions must be taken into account when explaining any product of radiative processes within them. Nevertheless an explanation in which there is an increase in the escape fraction of ionizing photons from \\HII\\ regions above this luminosity is still a reasonable way to account for a variety of observations, even though more sophisticated modelling will be needed to give final quantitative answers.Our sample has extended previous \\HII\\ region analyses, by including a significant number of faint \\HII\\ regions with their corresponding values of $\\sigma$, and this has allowed us to better pin down the luminosity at which a clean upper envelope in the $L-\\sigma$ relation emerges. Using the measured value as a secondary standard candle we find that the uncertainty in the position of the break point yields a contribution to the  uncertainty in the distance to M\\,83 (0.4 Mpc) which is not much greater than the uncertainty in the Cepheid distance itself (0.2 Mpc). As our observations cover less than 10\\% of the optical disk of M\\,83 this suggests that a more complete data set could be used to reduce this uncertainty."
    },
    "1005/1005.4965_arXiv.txt": {
        "abstract": "We report the detection of the CO $J=1-0$ emission line in three near-infrared selected star-forming galaxies at $z\\sim1.5$ with the Very Large Array (VLA) and the Green Bank telescope (GBT). These observations directly trace the bulk of molecular gas in these galaxies. We find H$_2$ gas masses of $8.3\\pm1.9\\times10^{10}\\ M_\\sun$, $5.6\\pm1.4\\times 10^{10}\\ M_\\sun$ and $1.23\\pm0.34\\times10^{11}\\ M_\\sun$ for BzK-4171, BzK-21000 and BzK-16000, respectively, assuming a conversion $\\alpha_\\mathrm{CO}=3.6\\ M_\\sun$ (K km s$^{-1}$ pc$^{2}$)$^{-1}$. We combined our observations with previous CO $2-1$ detections of these galaxies to study the properties of their molecular gas. We find brightness temperature ratios between the CO $2-1$ and CO $1-0$ emission lines of $0.80_{-0.22}^{+0.35}$, $1.22_{-0.36}^{+0.61}$ and $0.41_{-0.13}^{+0.23}$ for BzK-4171, BzK-21000 and BzK-16000, respectively. At the depth of our observations it is not possible to discern between thermodynamic equilibrium or sub-thermal excitation of the molecular gas at $J=2$. However, the low temperature ratio found for BzK-16000 suggests sub-thermal excitation of CO already at $J=2$. For BzK-21000, a Large Velocity Gradient model of its CO emission confirms previous results of the low-excitation of the molecular gas at $J=3$. From a stacked map of the  CO $1-0$ images, we measure a CO $2-1$ to CO $1-0$ brightness temperature ratio of $0.92_{-0.19}^{+0.28}$. This suggests that, on average, the gas in these galaxies is thermalized up to $J=2$, has star-formation efficiencies of $\\sim100$ $L_\\sun$ (K km s$^{-1}$ pc$^2$)$^{-1}$ and gas consumption timescales of $\\sim0.4$ Gyr, unlike SMGs and QSOs at high redshifts. ",
        "introduction": "\\begin{table*} \\centering \\caption{Summary of the VLA observations\\label{table:1}} \\begin{tabular}{lcccccrcc} \\hline Source & R.A.$^a$ & Dec.$^a$ & Frequency$^b$ & Bandwidth & Beam & Pos. Ang. & Cov. Fraction$^c$ & Rms$^d$ \\\\ & (J2000) & (J2000) & (GHz) & (MHz) &   &  & (\\%) &($\\mu$Jy/beam) \\\\ \\hline\\hline BzK-4171 & 12\\ 36\\ 26.516 & $+$62\\ 08\\ 35.35   & 46.760 & 100  & $1.69\\arcsec\\times1.56\\arcsec$ & $+09.8\\degr$  &   96       &  44 \\\\ &                &                    & 43.340 & 100 & $2.02\\arcsec\\times1.67\\arcsec$ & $-53.0\\degr$   &          & 62 \\\\ BzK-21000 & 12\\ 37\\ 20.597 & $+$62\\ 22\\ 34.60   & 45.710 & 100  & $1.79\\arcsec\\times1.61\\arcsec$ & $+18.2\\degr$ &   98       & 42 \\\\ &                &                    & 43.340 & 100 & $2.03\\arcsec\\times1.67\\arcsec$ & $-54.2\\degr$   &          &72 \\\\ BzK-16000 & 12\\ 36\\ 30.120 & $+$62\\ 14\\ 28.00   & 45.653 & $13\\times3.125$ & $1.96\\arcsec\\times1.60\\arcsec$  & $+48.9\\degr$ &  76    & 145\\\\ &                &                    & 45.635$^e$ & 100  & $1.90\\arcsec\\times1.62\\arcsec$ & $+44.3\\degr$ &        & 181 \\\\ \\hline \\end{tabular} \\begin{flushleft} \\noindent $^a$ VLA 1.4 GHz position from \\citet{Morrison2010}. \\noindent $^b$ Central frequency between the two 50 MHz IFs. \\noindent $^c$ Fraction of the CO line covered by the VLA observations. $^d$ Image noise level measured around the source position. $^e$ Combined continuum emission from two 50 MHz bandwidth channels at 45.585 GHz and 45.685 GHz.\\end{flushleft} \\end{table*}    \\begin{figure*} \\centering \\includegraphics[scale=0.38]{figure1a.ps} \\includegraphics[scale=0.38]{figure1b.ps} \\includegraphics[scale=0.38]{figure1c.ps} \\caption{VLA images of the CO $1-0$ emission line from our sources. Contour levels are: -2, 2, 3 and 4$\\sigma$. The crosses indicate the VLA 1.4 GHz position for our sources \\citep{Morrison2010}. \\label{fig:bzk_im}} \\end{figure*}    A remarkable step in our understanding of galaxy formation has come from the determination of the unobscured history of the star formation rate (SFR) density over a wide range of redshift. The SFR history of the Universe has a peak from $z=3$ to 1, and presents a steady decline at $z<1$ \\citep[e.g.][]{Lilly1996, Madau1996, Steidel1999}. Apparently, a significant contribution to this peak is provided by the vigorous bursts of star formation triggered by galaxy mergers and interactions. In these cases, the star formation is very efficient, showing high SFRs and high gas surface densities on scales of several kpc, similar to what is observed in distant submillimeter galaxies \\citep[SMGs;][]{Tacconi2006,Tacconi2008} and quasars \\citep[QSOs;][]{Walter2004, Walter2009, Riechers2008, Riechers2009}. However, an important fraction of the SFR density appears to come from normal star-forming disk galaxies \\citep{Bell2005, Elbaz2007, Genzel2008}. Here, continuous flows of cold gas from the intergalactic medium may provide the necessary fuel for star formation \\citep{Keres2005, Dekel2009}.  Studies of the properties of the molecular gas (e.g., CO) in galaxies during the main epoch of galaxy formation ($z=1-3$) have mostly focused on SMGs and QSOs due to the limited capabilities of millimeter and radio telescopes and instruments \\citep[e.g.,][]{Solomon2005, Weiss2005b, Riechers2006, Carilli2007, Weiss2007, Aravena2008, Coppin2008, Carilli2010}. These IR-luminous objects were found to have large reservoirs of molecular gas ($M(\\mathrm{H}_2)\\sim10^{10-11} M_\\sun$) that sustain SFRs of $\\sim500-1000\\ M_\\sun$ yr$^{-1}$ for $<100$ Myr. Similar to what is observed in local ultra-luminous infrared galaxies (ULIRGs), the molecular gas is in local thermodynamic equilibrium (LTE) up to high-$J$ CO transitions \\citep[$J>3$;][]{Riechers2006, Weiss2007}. This is expected for warm, dense H$_2$ molecular gas.  Recently, large amounts of molecular gas, similar to those observed in bright high-redshift SMGs and QSOs, were found in six relatively quiescent disk galaxies at $z\\sim1.5$ \\citep{Daddi2008, Daddi2009}, 14 similar galaxies at $z\\sim1.2$ and $z\\sim2$ \\citep{Tacconi2010} and three disk galaxies at $z=0.5$ \\citep[][; Salmi et al. in preparation]{Daddi2010b}. A detailed analysis of the properties of these near-IR selected galaxies revealed SFRs $\\sim50-200\\ M_\\sun$ yr$^{-1}$, with stellar masses of $\\sim10^{10-11}\\ M_\\sun$. These galaxies apparently have high molecular gas fractions, low star-formation efficiencies (SFEs) \\citep{Daddi2008, Daddi2009,Tacconi2010} and CO luminosity to gas mass conversion factors similar to that found in the Milky Way galaxy disk rather than that of starbursting ULIRGs and SMGs \\citep{Daddi2009}. This means that these objects have yet to convert a large fraction of their gas into stars, indicating galaxies right in the process of stellar build-up. The first CO $1-0$ measurements in two BzK galaxies were previously reported by \\citet{Dannerbauer2009}, but only led to a stacked detection at $\\sim3\\sigma$. Those observations, combined with the first detection of the CO $3-2$ line in one of these galaxies (BzK-21000), allowed \\citet{Dannerbauer2009} to perform an analysis of the excitation properties of the molecular gas, finding low excitation conditions, similar to what is seen in local disk galaxies. This suggests an important difference to what is found in more luminous objects at high-redshift, where the CO emission is in LTE up to high$-J$ CO transitions.  In this paper, we report individual detections of the CO $J=1-0$ emission line in a sample of three near-IR selected star-forming galaxies at $z\\sim1.5$ that have previously been studied in detail by \\citet{Daddi2009}: BzK-4171, BzK-21000 and BzK-16000. In section \\ref{section:obs} we describe our CO $1-0$ observations. In section \\ref{section:res}, we present our results, give fluxes and luminosities. In section \\ref{section:diss}, we summarize our results and discuss the average properties of the molecular gas in our galaxies. We assume a concordance cosmology with $H_0=71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_\\lambda=0.73$ and $\\Omega_m=0.27$. ",
        "conclusions": "\\label{section:diss}    \\begin{figure} \\centering \\includegraphics[scale=0.5]{figure6.ps} \\caption{Combined CO $1-0$ emission from the three BzK galaxies. Contours levels are: -2, 2, 3, 4, 5 and 6$\\sigma$. The cross indicates the reference VLA position.\\label{fig:stack}} \\end{figure}   We have detected the CO $1-0$ emission line in three massive star-forming galaxies at $z\\sim1.5$. Our observations allow us to carry out a direct comparison with studies of galaxies at low-redshift, and hence put unique constraints on the properties of the molecular gas in disk galaxies at high-redshift.  From our molecular gas excitation analysis for BzK-21000, we confirm the \\citet{Dannerbauer2009} results that the gas in this galaxy has low excitation conditions at $J=3$ (Fig. \\ref{fig:lvg}). The CO ladder for this galaxy seems to be thermalized up to $J=2$ and it is sub-thermal at $J=3$. The CO ladder is similar to that found for local disks, as in NGC6946 \\citep[Fig. \\ref{fig:lvg};][]{Mauersberger1999}. However, at the significance of this detection, the CO ladder is also consistent with the Milky Way galaxy, for which the emission appears to be non-LTE already at $J=2$.  A similar case is that of BzK-4171, where the relatively large uncertainty in the temperature ratio, $r_{21}$, does not allow us to differentiate between thermal or sub-thermal molecular gas up to $J=2$. Measurements of the CO $3-2$ transition are necessary to check whether or not higher order CO transitions are in LTE in this galaxy. Interestingly, in the case of BzK-16000, our measurements suggest that the CO emission is sub-thermal ($r_{21}<1$) already at $J=2$, resembling what appears to be the case for the inner disk of the Milky Way galaxy \\citep{Fixsen1999}. Our estimate of $r_{21}\\sim0.4$ differs from the LTE at the $\\approx2\\sigma$ level. Measurements of higher$-J$ transitions will help to study this particular galaxy in more detail.  Based on the previous CO $1-0$ observations of BzK-21000 \\citep{Dannerbauer2009}, \\citet{Daddi2009} and \\citet{Tacconi2010} assumed $r_{31}=0.5$ to convert their CO $3-2$ luminosities into CO $1-0$ luminosities. For this galaxy, we obtain $r_{31}\\sim0.6$, which validates, within $1\\sigma$, their assumption.  For a solid measurement of the average CO $1-0$ line emission and the average molecular gas properties of our BzK galaxies, we extracted $10\\arcsec\\times10\\arcsec$ cutouts centered at the VLA peak positions and produced a noise-weighted average CO $1-0$ map of their combined emission (Fig. \\ref{fig:stack}). We obtain a detection with a flux density of $225\\pm36\\ \\mu$Jy ($6.25\\sigma$). Stacking with respect to the CO position yields a flux density of $220\\pm40\\ \\mu$Jy (5.5$\\sigma$), consistent within the uncertainties ($<1\\sigma$) with the flux measured by stacking with respect to the VLA position. From the values presented in \\citet{Daddi2009} for the CO $2-1$ emission, we compute a noise-weighted average CO $2-1$ to CO $1-0$ brightness temperature ratio of $0.92_{-0.19}^{+0.28}$, suggesting that the CO SED from these galaxies is thermalized up to $J=2$. This ratio is consistent with what is found in local disk galaxies, which have lower CO luminosities than our BzK galaxies \\citep[][ and references therein]{Mauersberger1999, Bayet2006}, but it is also similar to other galaxy populations with CO luminosities comparable to that of our BzK galaxies, such as local ULIRGs \\citep{Weiss2005b, Guesten2006, Hitschfeld2008, Bayet2006, Bayet2009}, and high-redshift SMGs and QSOs \\citep{Solomon2005, Weiss2005b, Riechers2006, Weiss2007}. Thus, our comparison based only on CO $2-1$ to CO $1-0$ ratios is not yet sufficient to decide whether the excitation conditions in our BzK galaxies are, on average, different or not from those of other galaxy populations.  We note that \\citet{Daddi2009} assumed $r_{21}=0.86$ to convert the CO $2-1$ luminosities into CO $1-0$ luminosities for their sample of 6 BzK galaxies, following the \\citet{Dannerbauer2009} results on BzK-21000. This assumption is supported by our measurement from the average stacked CO map.  From our stacked CO $1-0$ measurement, we find on average a SFE $\\approx100$ (K km s$^{-1}$ pc$^2$)$^{-1}$. This, including individual values, is well below the average for high-redshift SMGs, $560\\pm210$ $L_\\sun$ (K km s$^{-1}$ pc$^2$)$^{-1}$ \\citep{Greve2005, Tacconi2006} and similar to the values found for local spiral galaxies \\citep{Boselli2002, Leroy2008} and for other disk galaxies at high-redshift \\citep{Tacconi2010}. However, it is marginally consistent with the average value found in local ULIRGs, $\\sim225$ $L_\\sun$ (K km s$^{-1}$ pc$^2$)$^{-1}$ \\citep{Solomon1997, Yao2003}. We obtain an average gas depletion lifetime of $\\sim0.4$ Gyr for the galaxies in our sample, which is higher than what is found in SMGs, $<0.05$ Gyr \\citep{Tacconi2006}.  Observations of disk galaxies at high-redshift (including this work) have provided detections only for the lower-$J$ CO transitions, and thus they probe the lower end of the CO SED. Recent studies of high-redshift SMGs and QSOs have indicated that two gas components can be present: a cold, low density gas component and a warm, dense component \\citep[e.g., ][]{Weiss2007, Carilli2010}. If indeed such is the case for star-forming BzK galaxies, observations of higher-$J$ CO lines are crucial. High-definition CO multi-transition studies over larger samples are necessary. Here, the EVLA and Atacama Large Millimeter Array (ALMA) will play a fundamental role, disentangling the molecular gas distribution of the cold molecular gas traced by CO $1-0$ and the warm gas probed by the higher-$J$ transitions."
    },
    "1005/1005.3890_arXiv.txt": {
        "abstract": "Interacting binaries consist of a secondary star which fills or is very close to filling its Roche lobe, resulting in accretion onto the primary star, which is often, but not always, a compact object. In many cases, the primary star, secondary star, and the accretion disk can all be significant sources of luminosity. SIM Lite will only measure the photocenter of an astrometric target, and thus determining the true astrometric orbits of such systems will be difficult. We have modified the Eclipsing Light Curve code \\citep{Orosz00} to allow us to model the flux-weighted reflex motions of interacting binaries, in a code we call \\textsc{reflux}. This code gives us sufficient flexibility to investigate nearly every configuration of interacting binary. We find that SIM Lite will be able to determine astrometric orbits for all sufficiently bright interacting binaries where the primary or secondary star dominates the luminosity. For systems where there are multiple components that comprise the spectrum in the optical bandpass accessible to SIM Lite, we find it is possible to obtain absolute masses for both components, although multi-wavelength photometry will be required to disentangle the multiple components. In all cases, SIM Lite will at least yield accurate inclinations, and provide valuable information that will allow us to begin to understand the complex evolution of mass-transferring binaries. It is critical that SIM Lite maintains a multi-wavelength capability to allow for the proper deconvolution of the astrometric orbits in multi-component systems. ",
        "introduction": "The field of stellar evolution has matured to the point that we generally understand the entire evolution of an isolated star if we know its initial mass. While many details remain to be solved, it is unlikely that SIM Lite will provide a revolution in our understanding of the evolutionary processes for single stars of low and intermediate mass. The same cannot be said for binary star systems, where there are at present a number of outstanding issues that remain unsolved. Among the most difficult is how mass transfer proceeds during the common envelope stage that occurs in the post-main sequence lives of close binaries. The formation of nearly every mass transferring binary requires a common envelope phase where the secondary star helps to strip the atmosphere of the primary while driving mass from the system, and shrinking the binary's orbit. Current theory \\citep[c.f.,][]{Pod03} has a difficult time explaining the creation of systems with massive black hole primaries (M$_{\\rm BH}$ $\\geq$ 9 M$_{\\sun}$) and very low mass secondary stars (M$_{\\rm 2}$ $\\leq$ 0.7 M$_{\\sun}$). Similar problems exist across the mass spectrum of close and interacting binaries (IBs), which for the purposes of the current study include systems with a black hole, neutron star, or white dwarf (WD) primary and a non-degenerate companion. For example, \\citet{Howell01b} explores the well known 2-3 hr period gap observed in cataclysmic variable (CV) systems, which is postulated to be a result of cessation of mass-transfer from the main-sequence, low-mass secondary star to the white dwarf companion for periods between 2-3 hours. \\citet{Howell01b} runs binary population synthesis models that show this theory is only correct if the secondary stars in these systems are up to 50\\% oversized, and thus both components are really much less massive than has been typically assumed via standard main-sequence mass-radius-temperature relations. Thus, if accurate masses can be determined for these systems, their formation theories can be directly tested.  Typically in binary star work, radial velocity curves are used to obtain masses for the components. While it can be quite simple to get a radial velocity curve for one component of an IB, (the secondary star if it dominates the luminosity, else typically the disk), it is difficult, or even impossible, to get such data for their black hole, neutron star, or WD primaries. Thus, along with the radial velocity amplitude of the secondary or disk, one must also assume a mass and know the binary orbital inclination in order to solve for the component masses. It is especially difficult to determine the binary inclination. In a large number of IBs the secondary star can be quite prominent at infrared wavelengths, and since the object is distorted, it exhibits ellipsoidal variations as it orbits the primary \\citep[c.f.][]{Gelino01}. The amplitude of these variations are dependent on the orbital inclination, and to a lesser extent, the properties of the secondary star. However, even a small amount of contamination by the accretion disk introduces considerable uncertainty in this method, and nearly all interacting binary systems have some level of contamination from the accretion process. Thus, this technique of using infrared ellipsoidal variations to determine the inclination fails if one cannot ascertain the spectrum and level of the contaminating source.  This paper will show how SIM Lite, with multi-wavelength, microarcsecond astrometry, will be the first mission capable of determining the multi-wavelength astrometric orbits of IBs. The field of optical astrometry has recently been making great progress from both ground-based observations, as well as is posed to make dramatic leaps with the upcoming space-based missions GAIA and SIM Lite. The limits of ground-based astrometry have been recently pushed by both CHARA and PRIMA/VLTI, but neither are capable of the kind of astrometric measurements needed to probe IBs. CHARA has multi-wavelength capabilities, but can only provide angular resolution to $\\sim$200 $\\mu$as \\citep{Brummelaar05}, which is far larger than the reflex motions to be discussed in this paper, and requires very bright targets. PRIMA/VLTI will achieve $\\sim$30-40 $\\mu$as precision \\citep{Belle2008}, which may be at the limit of usefulness for these systems, but only in K-band, which does not help distinguish individual component orbits, as will be shown. The GAIA mission will provide astrometry for $\\sim$10$^{9}$ objects with 4 - 160 $\\mu$as accuracy, for stars with V=10-20 respectively, and does posses multi-wavelength capabilities \\citep{Cacciari09}. However, GAIA is a scanning satellite and cannot perform pointed or time-critical observations, and in fact only achieves this accuracy by averaging $\\sim$80 individual measurements, the individual errors of which range from 36 - 1,431 $\\mu$as, again for stars with V=10-20 respectively \\citep{Mignard05}. Thus, GAIA cannot provide the high-precision, time-critical pointed observations needed to study IBs. In contrast, SIM Lite will be able to point at any desired object for any length of time, providing single-measurement accuracy of $\\sim$1 $\\mu$as \\citep{NASA09}. As well, SIM Lite will have $\\sim$80 spectral channels, spanning 450 to 900 nm, thus providing multi-wavelength, microarcsecond astrometry \\citep{NASA09}.  SIM Lite will offer the first chance at measuring astrometric orbits for a large number of IBs, thus directly yielding inclinations and allowing for the precise measurement of the masses of both components in these systems. Determining the inclination of a zero-eccentricity system from astrometry is as straightforward as determining the ratio the semi-major and semi-minor axes of the projected ellipse on the sky. Ideally, one can derive accurate values for the masses of both components from the astrometry of a single component \\citep{Benedict00}, but to do so requires that one knows the orbital period, the semi-major axis of the apparent orbit, the parallax of the system, and the mass ratio of the two components. SIM Lite can provide the first three of these four parameters, but the mass ratios must be estimated for many of the main systems of interest. However, for binaries where there is a significant amount of light from more than one component, \\emph{SIM Lite can determine absolute masses for both components}. Since SIM Lite only measures the location of the photocenter of the system, the observed reflex motion of the system will be wavelength-dependent. For example, in a binary where one component is hotter than the other, the measured motion in the blue part of the spectrum will be different from that measured in the red, with shorter wavelengths tracing the motion of the hotter component, and vice-versa. Thus, if one has at least two astrometric orbits at different wavelengths, and one knows the ratio of luminosities in each bandpass, (i.e. the spectra of the components), then one could reconstruct the individual orbits of each component and obtain absolute masses for both. Since SIM Lite is currently designed to have $\\sim$80 spectral channels, and since multi-color photometry and spectroscopy already exists for most IBs of interest, this technique should be quite feasible for most IBs.  In the next section we describe the procedure and code used to model the reflex motions of IBs. In section 3 we examine the results for ``proto-typical'' systems, including both X-ray binaries and cataclysmic variables, and estimate the required observing time required by SIM Lite to obtain reasonably accurate parameters for each system. We also briefly investigate how the presence of disk temperature gradients and hotspots in the systems affect the derived orbits. In section 4 we present a full modeling of simulated SIM Lite data for one system, and show the precision and accuracy of recovered astrometric parameters and derived system masses. We summarize our results in section 5. ",
        "conclusions": "We have presented a modeling code, \\textsc{reflux}, that is capable of modeling the multi-wavelength astrometric signature of interacting binaries. We have presented models for multiple IB systems, and find that for primary or secondary star dominated systems, the contamination of the photocenter is minimal and SIM Lite will provide good astrometric orbits if the system is bright enough, yielding absolute inclinations and indirect masses. For mixed component systems, SIM Lite will be able to directly determine absolute masses for both components when multi-wavelength astrometric curves are obtained and combined with existing spectroscopy and multi-color photometry. For disk-dominated systems, SIM Lite will only be able to obtain astrometric orbits for a few systems, due to the small relative motion of the compact object that the disk surrounds, but these should yield accurate inclinations. We find that while multi-wavelength SIM Lite data will be unable to distinguish between various disk temperature gradients and disk hotspots, it should be able to determine the location of magnetic threading regions in polars, and thus similar effects in other systems, via narrow wavelength astrometry. In total, SIM Lite should contribute greatly to our understanding of interacting binary systems, especially if its multi-wavelength capability is maintained in the final design."
    },
    "1005/1005.4681_arXiv.txt": {
        "abstract": "We present deep $J$ and $K_{s}$ band photometry of 20 high redshift galaxy clusters between $z=0.8-1.5$, 19 of which are observed with the MOIRCS instrument on the Subaru Telescope. By using near-infrared light as a proxy for stellar mass we find the surprising result that the average stellar mass of Brightest Cluster Galaxies (BCGs) has remained constant at $\\sim9\\times10^{11}M_{\\odot}$ since $z\\sim1.5$. We investigate the effect on this result of differing star formation histories generated by three well known and independent stellar population codes and find it to be robust for reasonable, physically motivated choices of age and metallicity. By performing Monte Carlo simulations we find that the result is unaffected by any correlation between BCG mass and cluster mass in either the observed or model clusters. The large stellar masses imply that the assemblage of these galaxies took place at the same time as the initial burst of star formation. This result leads us to conclude that dry merging has had little effect on the average stellar mass of BCGs over the last $9-10$ Gyr in stark contrast to the predictions of semi-analytic models, based on the hierarchical merging of dark matter haloes, which predict a more protracted mass build up over a Hubble time. We discuss however that there is potential for reconciliation between observation and theory if there is a significant growth of material in the intracluster light over the same period. ",
        "introduction": "Brightest Cluster Galaxies (BCGs) are the most luminous objects in the Universe in terms of stellar light and appear to be a separate population from bright ellipticals \\citep{sandage1972,sandage1976,bhavsar1985,oegerle1991,postman1995,bernstein2001,bernardi2007,vonderlinden2007,vl08,lin2009}. They reside in the deep potential wells of the cores of rich galaxy clusters, thought to descend from the first regions where mass began to accumulate after the Big Bang. Their luminosities and unique environments make them ideal candidates for the testing of the mass build up in galaxies across a large fraction of the Hubble time.  A number of studies have attempted to constrain the formation epoch and evolution of BCGs by comparing their $K$ band Hubble diagram to a range of stellar population models (e.g. \\citealt{abk98}; ~\\citealt{cm98}; ~\\citealt{nel02}). BCGs have been shown to follow passive evolution out to moderate redshifts but then differing results are seen at $z>0.5$, with some groups seeing a continuation of the passive trend (e.g. \\citealt{cm98}) while others claim the high redshift BCGs are fainter and therefore less massive than their local counterparts (e.g. \\citealt{abk98}). This is explained by a dependence of BCG luminosity on the mass of its host cluster \\citep{bcm00}, with most studies now agreeing that the evolution of BCGs in massive clusters can be described as passive to $z\\sim0.8$ \\citep{stott08,whey08}. However there is still some debate over the recent merging activity of BCGs particularly in lower mass clusters and Brightest Group Galaxies with a number of studies identifying major merger candidates \\citep{mulchaey2006,rines2007,tran2008}.  The favoured model for galaxy formation and evolution is via the hierarchical merging of dark matter haloes (e.g. \\citealt{davis85}).  In this model the mass of a galaxy gradually increases as it merges with neighbouring systems. A major development in the field has been the advent of large cosmological simulations such as the Millennium N-body Simulation \\citep{springel05} which models this hierarchical mass build up of dark matter haloes in a co-moving box 500\\,$h^{-1}$\\,Mpc (where $h=H_0/100\\,{\\rm km\\, s^{-1}\\,Mpc^{-1}}$) on the side.  Semi-analytic models are commonly used to describe the complex baryonic physics of galaxies in the context of the merger histories of dark matter halos within N-body simulations (e.g. \\citealt{b06}; \\citealt{deluc07}). These models are an efficient way to describe the competing processes affecting baryonic matter such as those that trigger or suppress star formation i.e. gas rich galaxy merging and feedback from active galactic nuclei or supernovae. The output of observables, such as galaxy magnitudes, from the semi-analytic models has proved to be valuable for astronomers as the modelled systems can be compared directly with measurable quantities. For this reason semi-analytic models can be a useful tool for making mock catalogues, predicting the behaviour of galaxies, testing cosmological theory and assessing the feasibility of telescope observations.  One important advantage of using BCGs to study galaxy evolution is that their theoretical counterparts can be easily and unambiguously identified as the central massive galaxies in mock clusters at the same redshift as the real systems.  \\cite{deluc07} presented the evolution of BCGs over a Hubble time, in a semi-analytic model based on the Millennium Simulation. This paper predicted that the stellar population of BCGs forms early, with 50\\% of the stellar mass in place by $z=5$ and 80\\% by $z=3$. However this early star formation takes place in separate components that gradually assemble into the BCGs seen in the local Universe, mainly through dry mergers that do not trigger additional star formation. So for example, at $z=1.5$ the sum of the stellar mass in all sub-components is over 90\\% of the mass of the fully assembled BCG at $z=0$, whereas the stellar mass in the main progenitor is on average only 20\\% of the galaxy's final mass.  From an observational standpoint, a significant number of high redshift galaxy clusters ($z\\gtrsim1$) have been discovered in recent years with X-ray surveys (e.g. \\citealt{rosati2004}; \\citealt{mullis05}; \\citealt{stan06}; \\citealt{brem06}; \\citealt{lamer08}). The {\\it XMM} Cluster Survey (XCS, \\citealt{romer01}; \\citealt{sh08}) is one such project performing a serendipitous survey to discover clusters in the XMM-Newton archive. The main goals of the XCS are to constrain cosmological parameters, measure the evolution of the hot gas through analysis of the cluster scaling relations and study galaxy evolution in clusters since the high mass cluster cores are thought to be environments comparable across all epochs. Furthermore X-ray luminosity and temperature are excellent proxies of cluster mass and enable us to directly compare the properties of real cluster galaxies with mock galaxies in similar cluster halo environments. With the advent of large wide-field optical and infrared surveys it is also possible to photometrically select high redshift galaxy clusters  based on the properties of their constituent galaxies, further increasing the number of known $z>1$ clusters (e.g. \\citealt{RCS04}; \\citealt{swin07}; \\citealt{Eisen08}).  In this paper we present near-infrared observations of 20 high redshift galaxy clusters ($0.8<z<1.5$), derive stellar masses for their BCGs  and compare to the latest semi-analytic models based on the Millennium N-body Simulation \\citep{deluc07,springel05}. We test how robust our results are to different star formation histories produced by three independent stellar population codes: \\cite{bc03}; \\cite{mar2005} and BaSTI \\citep{piet2004,percival2009}. This work builds on the result of \\cite{collins09} in which we found that the average stellar mass of the BCGs in 5 of the highest redshift galaxy clusters is not significantly different to that of BCGs in the local Universe.  Unless otherwise stated a Lambda Cold Dark Matter ($\\Lambda$CDM) cosmology ($\\Omega_{M}=$0.3, $\\Omega_{Vac}=$0.7, $H_{0}=$70 km s$^{-1}$ Mpc$^{-1}$) is used throughout this work. ",
        "conclusions": "We have demonstrated that the average stellar mass of BCGs in the highest redshift X-ray clusters is discrepant with those from similar mass dark matter halos in semi-analytic models based on the Millennium Simulation. Instead of the gradual build up of mass through dry merging predicted by \\cite{deluc07}, our observations suggest that the stellar mass in these objects has remained unchanged over the last 9 - 10 Gyr requiring a more rapid build up of stellar mass before $z=1.5$ some $4-5$ Gyrs after the big bang.  The timescale for the mass assemblage is similar to the age of the component stars ($2-3$~Gyrs), a situation that appears to resemble classical monolithic collapse \\citep{ebs62, lar74} rather than hierarchical formation. To form a galaxy of stellar mass $10^{12}\\,$ M$_{\\odot}$ over 4 Gyrs requires a mass deposition rate of about 250 \\,M$_{\\odot}$ yr$^{-1}$ and an efficient mechanism to feed the gas into the inner regions of the halo where it can form stars. Unfortunately the merging process becomes inefficient for massive galaxies because merger induced shocks lead to heating as opposed to radiative cooling of the gas \\citep{binney04}. One suggestion is that the early assembly of massive galaxies at $z\\geq 2$ is driven by narrow streams of dense cold gas which penetrate the shock-heated region greatly increasing the efficiency of the gas deposition and associated star formation \\citep{birnboim2003, keres2005, dekel09}.  Alternatively, a deficiency may lie in the semi-analytic treatment of the physical processes in the densest environments during early hierarchical assembly; this contention is supported by the fact that current predictions are moderately consistent with observations of the evolution of luminous red galaxies \\citep{wake06, alm08}, whereas our results, which focus on the most massive subset of this population, the BCGs, differ much more from the model predictions.  From a theoretical point of view this area of research is constantly evolving due to the challenges posed by observation and the increased availability of powerful computers. A recent high resolution simulation of a single cluster \\citep{rands09} predicts that 50\\% of the BCG final mass in a massive (10$^{15}$M$_{\\odot}$) cluster has built up by hierarchical merging in the last 9 billion years; decreasing the discrepancy between our findings and theory. As this is for one cluster and not a full cosmological simulation on the scale of the Millennium Simulation it is difficult to know whether this improvement is due to low number statistics or a better approximation of N-body and stellar-dynamical effects.  One consequence of our result is that if merging is important at all since $z\\simeq1$, the evolution of BCGs must be dominated by minor rather than major mergers, since the stellar mass appears unchanged since this time. Our observations are broadly consistent with the relatively low number of dry major mergers found at low redshift \\citep{liu09} and the model predictions of \\cite{khoch09} which show that only $10-20\\%$ of galaxies more massive than $6.3\\times10^{10}\\, M_{\\odot}$ have experienced dry major mergers within their last Gyr at any given redshift $z\\leq1$. Numerical simulations find that the scale sizes of early galaxies can grow from dry minor merging by a factor $2-3$ since $z=1$, e.g. \\cite{naab09}  and suggest that the cD--like haloes of BCGs are formed late, resulting in the relatively recent departure of BCGs from the Kormendy relation for ordinary elliptical galaxies \\citep{rands09}.  One possibility which might help reconcile observations and theory is the growth of stellar material constituting the diffuse intracluster light within the cluster cores. From dissipationless simulations of dark matter haloes, \\cite{conroy2007} find that while BCGs do not appear to evolve strongly at $z<1$, the intracluster light surrounding such galaxies is growing substantially, with up to $\\sim 85\\%$ of the stars in the intracluster medium of present day clusters deposited at $z<1$ (see also \\cite{willman2004}; \\cite{murante2007}; \\cite{Conroy2007b}; \\cite{White2007}; \\cite{Henriques2010}). This inside-out growth is broadly consistent with the dry minor merging scenario for the local elliptical galaxy population (\\citealt{bezanson2009}; \\citealt{hopkins2010}) required to explain their rapid size increase since $z\\sim2$ \\citep{vandokkum2009,vandokkum2010}.  Observationally, recent results have confirmed the overall importance of  intracluster light: \\cite{rudick2006} find that the intracluster light constitutes $\\sim 10\\%$ of the entire  cluster stellar light; a result that appears to hold even for non-cD clusters \\citep{feldmeier2004}. From surface photometry out to 300 kpc of  24 clusters at $z<0.1$, \\cite{gonzalez2004} demonstrate that the outer cD component of BCGs traces the cluster potential and has $\\sim 10$ times the total luminosity of  the inner elliptical profile. These results suggest that further work on the growth of stellar light in the outer parts of BCGs is required to provide a consensus."
    },
    "1005/1005.5291_arXiv.txt": {
        "abstract": "s{First order phase transitions in the early universe can give rise to a stochastic background of gravitational waves. A hypothetical first order electroweak phase transition is particularly interesting in this respect, since the signal is in the good frequency range to be detectable by the space interferometer LISA. Three main processes lead to the production of the gravitational wave signal: the collision of the broken phase bubbles, the magnetohydrodynamical turbulence in the plasma stirred by the bubble collisions, and the magnetic fields amplified by the magnetohydrodynamical turbulence. The main features of the gravitational wave spectrum, such as the peak frequency, the amplitude, and the slopes both at low and high wave-number can be predicted by general arguments based on the characteristics of the source: in particular, the structure of its space and time correlation. We find that the gravitational wave signal from a first order phase transition occurring at electroweak symmetry breaking falls into the LISA sensitivity range if the phase transition lasts for about one hundredth of the Hubble time and the energy density of the turbulent motions is about twenty percent of the total energy density in the universe at the phase transition time.} ",
        "introduction": "It is likely that in the years to come the interferometers LIGO \\cite{LIGO} and VIRGO \\cite{VIRGO} will provide the first direct detection of gravitational waves (GWs). Terrestrial interferometers operate in the frequency range 10-1000 Hz, while the space-based interferometer LISA \\cite{LISA} will have its best sensitivity around the milliHertz frequency. A possible target of these experiments is a stochastic background of GWs of cosmological origin \\cite{maggiore}. The detection of such a background, relic of the early universe, would have a profound impact both on cosmology and on high energy physics. Once GWs have been emitted, in fact, they propagate freely through the universe, carrying direct information on the physical process that generated them: their detection would therefore provide a new way of probing the primordial universe, and correspondingly physics at very high energies, which would not be accessible otherwise.  One of the processes that can generate such a stochastic background of GWs is a relativistic first order phase transition \\cite{hogan83,witten,hogan86,TW,beta}. In the course of its adiabatic expansion, the universe might have undergone several phase transitions driven by the temperature decrease. The nature of the phase transitions depends on the particle theory model, but if they are first order they proceed through the nucleation of broken phase bubbles, which is a very violent and inhomogeneous process capable of sourcing GWs \\cite{Kos92,Kos93,Kam94}. In the following we review the different mechanisms of generation of GWs by a first order phase transition, and show that the characteristic frequency and amplitude of the GW signal are related respectively to the temperature at which the phase transition occurs and to its strength. We also review the analytical method developed in Refs.~\\cite{bubble,thomas,MHD}, which predicts the main features of the GW spectrum starting from a description of the source which we tried to make as simple and as model independent as possible. The GW signal depends explicitly on a few free parameters: under specific choices for these parameters, it can be potentially interesting for observation with LISA \\cite{apreda,nicolis,kahn1} or, more speculatively, advanced LIGO. ",
        "conclusions": ""
    },
    "1005/1005.2687_arXiv.txt": {
        "abstract": "*{ The local universe is the best known part of our universe. Within the CLUES project ({\\tt http://clues-project.org} --- Constrained Local UniversE Simulations) we perform numerical simulations of the evolution of the local universe. For these simulations we construct initial conditions based on observational data of the galaxy distribution in the local universe. Here we review the technique of these constrained simulations. In the second part we summarize our predictions of a possible Warm Dark Matter cosmology for the observed local distribution of galaxies and the local spectrum of mini-voids as well as a study of the satellite dynamics in a simulated Local Group.  }   \\abstract{ The local universe is the best known part of our universe. Within the CLUES project ({\\tt http://clues-project.org} --- Constrained Local UniversE Simulations) we perform numerical simulations of the evolution of the local universe. For these simulations we construct initial conditions based on observational data of the galaxy distribution in the local universe. Here we review the technique of these constrained simulations. In the second part we summarize our predictions of a possible Warm Dark Matter cosmology for the observed local distribution of galaxies and the local spectrum of mini-voids as well as a study of the satellite dynamics in a simulated Local Group.  } ",
        "introduction": "\\label{sec:intro_h009z}   During the last decade cosmological parameters have been determined to a precision of just a few percent giving rise to the standard model of cosmology: a flat Friedmann universe whose mass-energy content is dominated by a cosmological constant (the $\\Lambda$ term), a cold dark matter (CDM) component and baryons. The convergence to a standard model of cosmology sets the framework for studying the formation of both the large scale structure and of galaxies within that model. The main goal of this study is to achieve a physical understanding of how the observed structure emerged within the context of the given cosmological model.  This model specifies the cosmological expansion history, the initial conditions and the material content of the universe. These plus the knowledge of the physical laws governing the dynamics of the dark matter, baryons and radiation provide the framework within which a successful model of galaxy formation can be developed.  The basic paradigm of structure formation in the $\\Lambda$CDM was formulated more than 30 years ago\\cite{whiterees78}.  It suggests that dark matter (DM) clusters to form DM halos, within which galaxy formation takes place via complex baryonic physics.  Intensive numerical efforts of the last decade have validated the basic White \\& Rees picture.  Yet, the process of galaxy formation is much more complicated than that simple picture.  DM halos grows via process of mergers of substructures and galaxy formation proceeds by the combined action of clumpy and anisotropic gas accretion and mergers dwarf galaxies. As dwarf galaxies cross the virial radius of the DM halo they become satellites, subject to tidal forces exerted by the central potential.  This transforms the satellites into tidal streams and eventually phase mixing dissolves the streams into the stellar halo of galaxies. This process manifests itself in the combined color and dynamical phase space of galactic stella halos.  This is the universal mode of galaxy formation but it can be observed and analyzed only in the very local universe, resulting in the so-called near-field cosmology. This motivates cosmologists to turn their attention and study the archeology of the Local Group (LG) in their quest for understanding galaxy formation. This also motivates us to simulate the formation of the LG in the most realistic possible way.  Cosmological simulations may cover a large dynamical and mass range. A representative volume of the universe should be large, but this comes at the expense of the mass resolution, which must be decreased in proportion to the simulated volume size.  One may reduce the box size at the expense of possibly not being representative. To overcome this problem we use smaller simulation boxes specifically designed to represent the observed local universe.  The algorithm of constrained realizations of Gaussian field provides a very attractive method of imposing observational data as constraints on the initial conditions and thereby yielding structures which can closely mimic those in the actual universe.  The prime motivation of the CLUES project is to construct simulations that reproduce the local cosmic web and its key 'players', such as the Local Supercluster, the Virgo cluster, the Coma cluster, the Great Attractor and the Perseus-Pisces supercluster.  The main drawback of current constrained simulations is that they do not directly constrain the sub-megaparsec scale structure, yet they enable the simulation of these scales within the correct environment.  Such simulations provide a very attractive possibility of simulating objects with properties close to the observed LG and situated within the correct environment.  Their random origin, however, limits the predictive power of the simulations and in particular the constraining power of our 'near field cosmology' experiments. It is one of our main goals to improve the constraining power and ability of the simulations to reproduce the actual observed LG. ",
        "conclusions": "\\label{sec:summary_h009z}  Within the CLUES project ({\\tt http://clues-project.org}) we have performed a series of constrained simulations of the local universe. These simulations reproduce the observed large scale structures around the Local Group whereas small scale structures below scales of $\\sim 1 \\hMpc$ are essentially random. In four simulations we have identified Local Groups situated at the right position but having different merging histories. Constrained simulations are a useful tool to study the formation and evolution of our Local Group in the right cosmological environment and the best possible way to make a direct comparison between numerical results and observations, minimizing the effects of the cosmic variance.    \\begin{acknowledgement} The computer simulations described here have been performed at LRZ Munich and BSC Barcelona.  We acknowledge support by ASTROSIM, DFG, DEISA and MICINN (Spain) for our collaboration.  We thank Jesus Zavala (Munich), Noam Libeskind (Potsdam), Jarek Klimentowski (Warsaw) and Kristin Riebe (Potsdam) for providing plots from our common research projects.    \\end{acknowledgement}"
    },
    "1005/1005.2825_arXiv.txt": {
        "abstract": "We present EZ (Easy redshift), a tool we have developed within the VVDS project to help in redshift measurement from otpical spectra. EZ has been designed with large spectroscopic surveys in mind, and in its development particular care has been given to the reliability of the results obtained in an automatic and unsupervised mode. Nevertheless, the possibility of running it interactively has been preserved, and a graphical user interface for results inspection has been designed. EZ has been successfully used within the VVDS project, as well as the zCosmos one. In this paper we describe its architecture and the algorithms used, and evaluate its performances both on simulated and real data. EZ is an open source program, freely downloadable from http://cosmos.iasf-milano.inaf.it/pandora. ",
        "introduction": "Thanks to larger telescopes and more powerful instruments, during the last decade we have witnessed an explosion in the size of spectroscopic surveys both in the nearby and in the more distant Universe:  from the hundreds of objects of the surveys carried out in the early nineties (ESP, \\citet{ESP},CFRS, \\citet{CFRS}) to the thousands of galaxies of present days deep spectroscopic programs (zCosmos, \\citet{zcosmos_main}, VVDS, \\citet{vvds_main,vvds_wide} , DEEP2, \\citet{deep2}) and to the hundred thousands of more nearby surveys (e.g. 2dF, \\citet{2df}, SDSS, \\citet{sdss7}, 6dF, \\citet{6df}) or to the hundred thousand or million objects of the just started programs (VIPERS, \\citet{VIPERS}, WiggleZ, \\citet{WiggleZ}, BOSS, \\citet{BOSS}) and the few millions of galaxies of the future (e.g. EUCLID, \\citet{euclid}).\\\\ Such an evolution in survey size requires an analogous step forward in the way the data are reduced and analyzed, making compulsory the use of reliable automatic tools. Among the various tasks to be done when carrying out a redshift survey, redshift measurement is one of the most demanding in terms of human resources and skills required. Historically, the {\\it rvsao} \\citep{rvsao} package within IRAF has been the first widely used program which has been adopted for redshift measurement. It makes use of emission lines ({\\it emsao} task) and of the continuum shape ({\\it xcsao} task), on the basis of a correlation of the spectral continuum against galaxy or stellar templates, implementing the algorithm originally proposed by \\citet{tonryDavis}. Although very powerful, it does not foresee to weigh spectral features by an error spectrum. When the VIMOS-VLT Deep Survey (from here on, VVDS) started, it was immediately understood that this could be a major drawback in using {\\it rvsao} within that project: VIMOS Low Resolution Red (LR-Red) spectra, like those obtained within the VVDS project, are affected by heavy fringing redwards of $\\sim8000\\mbox{\\AA}$. As a consequence, spectra show features which are purely due to fringing but which can easily be mistaken with emission lines by any automatism, as well as by humans, unless they are properly weighted. Furthermore, if both emission and absorption lines are to be used, some scripting is required to couple together results from the pure cross correlation algorithm and the emission lines search one. Last, but not least, the IRAF implementation of {\\it rvsao} did not allow a direct interface with the VIPGI package used for VVDS Data reduction \\citep{vipgi}.\\\\ For all these reasons, we have decided to develop EZ (Easy redshift, pronounced ``easy''), a new tool which inherits from the {\\it rvsao} experience and the cross correlation technique, but allows to be easily plugged within VIPGI and could also be used as a VO-compliant tool via SSAP \\citep{SSAP} or other communication protocols. From the user's point of view, the most important requirement was to be able to run the program both in a completely unsupervised and in fully interactive mode, with a user friendly graphical interface to inspect the results obtained using different parameters or templates. Furthermore, it must be possible to use templates which do not necessarily cover the full spectral range for any redshift value. Finally, it has to be easily usable for different projects, i.e. all parameters must be defined either through a project dependent parameter file or via command line. EZ has been initially based upon the prototypes  KBred (Scaramella 1999, unpublished) and YAZ (Scaramella 2004, VVDS team Internal report), with which the first 10000 VVDS spectra were measured, and which contained the correlation and fitting algorithm described below. \\\\ In section \\ref{concept} we illustrate the rationale of EZ. Section \\ref{architecture} gives a short illustration of the software architecture and the most important algorithms used are illustrated in section \\ref{algorithm}. Section \\ref{dec_tree_verbose} describes in full how a redshift can be measured by EZ and section \\ref{reliability} illustrates how a reliability flag is attached to the best solution proposed. In sections \\ref{results_sim} and \\ref{results_real} we evaluate the performances on simulated and real data respectively. Finally section \\ref{gain} shows how much time and efforts can be spared using a tool like EZ when carrying out a large redshift survey. ",
        "conclusions": "We have presented, EZ, an automated tool devoted to redshift measurement. The concept at the basis of the tool is the {\\it decisional tree}, i.e. the sequence of operations to be performed to obtain a redshift. Such sequence of operations can be customized according to the kind of sample at hand: e.g. by discarding any check on stars, or using only emission lines without performing the correlations on the continua in the case of spectra obtained with an instrument operating in slitless mode. \\\\ EZ can be used both interactively, with the help of a graphical user interface, or totally blindly in unsupervised mode. It is developed in Python, with the bulk of computations performed in C to increase its computational speed. Its Python classes can be directly imported in any other Python based program, thus making it fully embeddable in any application.\\\\ Within EZ we have developed a method to assign a reliability flag on the measurement obtained, in order to mimic the kind of reasoning done by astronomers when assessing the goodness of the solution found. The implemented flagging system, though, is rather conservative.\\\\ We have tested EZ on VIMOS-like simulated data, and have shown that its performances are exceptionally good, the redshift measurement success rate being above 95\\%. \\\\ We have blindly applied EZ to  the VVDS-Deep, VVDS-Wide and zCosmos bright samples, and have demonstrated how this tool behaves very well also on real data. The success rate obtained is around 70\\%, and rises above 90\\% for redshifts classified as {\\it very secure} by astronomers.\\\\ Finally, we have shown that the adoption of a similar tool can save from a minimum of 50\\% to a maximum of 95\\% of the redshift measurement load, according to the quality of the data and to the degree of completeness and purity of results one is willing to sacrifice in favor of a  fast output. \\\\ EZ is now being routinely and blindly run as part of the reduction process within the VIPERS survey, while a customized version has been set up to perform simulations of the E-NIS spectrograph \\citep{euclid}.\\\\ EZ is an open source program, freely downloadable from http://cosmos.iasf-milano.inaf.it/pandora"
    },
    "1005/1005.5402_arXiv.txt": {
        "abstract": "It is well established that both total and spectral solar irradiance are modulated by variable magnetic activity on the solar surface. However, there is still disagreement about the contribution of individual solar features for changes in the solar output, in particular over decadal time scales. Ionized \\caii K line spectroheliograms are one of the major resources for these long-term trend studies, mainly because such measurements have been available now for more than 100 years. In this paper we introduce a new \\caii K plage and active network index time series derived from the digitization of almost 40,000 photographic solar images that were obtained at the 60-foot solar tower, between 1915 and 1985, as a part of the monitoring program of the Mount Wilson Observatory. We describe here the procedure we applied to calibrate the images and the properties of our new defined index, which is strongly correlated to the  average fractional area of the visible solar disk occupied by plages and active network. We show that the long-term variation of this index is in an excellent agreement with the 11-year solar cycle trend determined from the annual international sunspot numbers series. Our time series agrees also very well with similar indicators derived from a different reduction of the same data base and other \\caii K spectroheliograms long-term synoptic programs, such as those at Kodaikanal Observatory (India), and at the National Solar Observatory at Sacramento Peak (USA). Finally, we show that using appropriate proxies it is possible to extend this time series up to date, making this data set one of the longest \\caii K index series currently available. ",
        "introduction": "The importance of \\caii K spectroheliogram time series for the study of solar magnetism and irradiance variability is well established. Observations of the solar surface reveal magnetic fields with complex hierarchical structures, evolving on a wide range of different spatial and temporal scales. The most prominent aspect of this variability is the solar cycle of activity, with a period of approximately 11 years for the sunspot numbers and a period of about 22 years for the magnetic polarity. Thirty years of satellite measurements of the Sun's energy output have revealed that also the solar irradiance changes over the full range of time scales from minutes to decades ({\\it e.g.} \\opencite{2004A&ARv..12..273F}), and this variability is wavelength dependent \\cite{2001GeoRL..28.4119L,1998A&A...329..747S}. Empirical models have shown that the variability in solar irradiance is indeed modulated by the area variations of the solar surface magnetic features, to a high degree of correlation \\cite{2003AGUFMSH31C..01S,2004ApJ...611L..57F,1988ApJ...328..347F}. Two of these features, plages and chromospheric magnetic network, account for a significant portion of the Sun's total magnetic flux, UV and EUV variability, that play a critical role in determining the conditions in the heliosphere which directly influence the Earth's magnetosphere (see the references in the paper by \\inlinecite{2009SoPh..255..229F} for a discussion on this subject).  Observations near the core of the ionized calcium K line (393.37nm) are one of the most effective tools to investigate the morphology and evolution of both plages and chromospheric magnetic network. These measurements have been available since the early years of the twentieth century and, because of their correlation to the solar irradiance, they have been widely used as proxies to reconstruct the history of the solar magnetism and solar irradiance over the last 100 years and beyond \\cite{2003EOSTr..84..205F,2004ApJ...611L..57F}. The Mount Wilson archive of ionized \\caii K line spectroheliograms provides a fundamental resource for these studies. The intensity calibration of these images, however,  is a very difficult task because the disk brightness and level of background vary from one image to another due to contributions from the quiet Sun limb darkening curve of the Ca K-line, the geometrical distortion introduced by the guider errors that alter the image shape, and the vignetting function produced by the misalignment of the optical axis with respect to the center of the grating. For those observations that include the step wedge densitometer strips it is possible to derive a characteristic curve (also known as Hurter--Driffield curve, or H\\&D curve) to interpret the transparency of the photographic plate material in terms of an exposure quantity \\cite{1968ApOpt...7.1513D}. Unfortunately, the step wedge exposures only began 9 October 1961 making this approach unsuitable for the calibration of the entire database. However, digital filters are convenient tools for extracting the properties of the spatial intensity distribution of an image. In this paper we use a median (low-pass) filter to determine the image background of each \\caii K observation. A flat-fielded image is then obtained by dividing the image by its background. Two of the parameters that describe the pixel distribution of this normalized image can be used to define a \\caii K plage index and a plage contrast time series. The very good agreement between these two time series and similar products validates the analysis described in this paper \\cite{2009SoPh..255..229F,2009SoPh..255..239T}. Finally, we show here that it is possible to extend the length of the plage index time series up to date by using its excellent correlation to the Mount Wilson magnetic plage strength index (MPSI), making this \\caii K index times series one of the longest currently available. ",
        "conclusions": "In this paper we have described a method to study the time dependence of solar surface magnetism from the analysis of more than 70 years of \\caii K spectroheliograms obtained at the 60-foot solar tower in Mount Wilson between 1915 and 1985. These measurements are important because \\caii K emission is closely related to the magnetic field \\cite{1998SoPh..177..265J} and can provide a proxy for long-term induced changes, solar-cycle and beyond, in the total and spectral solar irradiance ({\\it e.g.} \\opencite{2000SSRv...94..139F}). Our method is based on the photometric properties of each individual solar images, using parameters computed from a histogram analysis of their intensity distribution. An important feature of this approach is the fact that it can be accomplished in a fully automated mode, without relying on a visual inspection of the images. We found that two of the parameters defining this intensity distribution, the constant baseline and the width, are very well correlated to the fractional area of the solar disk covered by plages plus active network and the plage contrast as computed by an independent investigation of the same database. The constant baseline measures the strength of the widely distributed regions that are brighter than average in \\caii K, and it is defined in this work as a CaK index. This index shows a temporal behavior which is in excellent agreement with the cycle of solar activity described by the international sunspot number. This is clearly illustrated in Figure \\ref{cak}, where the relative strength of solar cycles 15 to 21 is well reproduced.  Another remarkable property of this index is its correlation with the MPSI measurements shown in Figure \\ref{mpsi}. The relationship between these two variables is mostly linear, up to a MPSI value of about 2. Above this value the correlation drops quite significantly. Those are measurements corresponding to the time interval from 1979 to 1982.4, around the maximum of solar cycle 21, where both the MPSI and the CaK index show no systematic trends. As illustrated at the bottom of Figure \\ref{mpsi}, our adopted model produces an excellent agreement between these two quantities during both the period of minimum of activity, around 1975, and the maximum of cycle 21. The validity of this model, however, is limited to MPSI values up to about 7. In fact, above this limit, the asymptotic behavior of the model significantly reduces the correlation between these two indices of solar activity. Fortunately, an investigation of the historical record of MPSI daily measurements shows that this event is extremely rare, and will not affect the conclusions reached in this paper.  The possibility of extending the CaK index series beyond the time for which the \\caii K spectroheliograms are available, using the MPSI measurements as a proxy, is one of the main results of this work. This unique extended series, that covers almost a century of measurements, is shown in Figure \\ref{xcak} and tabulated in Table 1. Because of the close connection among \\caii K emission, solar surface magnetic field, and Total Solar Irradiance (TSI), this CaK time series can provide a proxy for studies on solar-cycle induced changes in the TSI that may play a role in climate change \\cite{2002AdSpR..29.1933S,2000SSRv...94..337S}. In particular, the progressive increase in the strength of the solar cycle of activity between 1928 and 1958 indicated by the sunspot number is closely tracked by the increase in strength of solar UV radiation as indicated by our CaK index. This probable increase in the solar output during a period of increasing greenhouse gas abundance needs to be included in evaluations of the global warming process.  \\begin{figure} \\begin{center} \\includegraphics[width=1.0\\textwidth]{mgii.eps} \\caption{Comparison between the CaK index (blue points) calculated in this work and the Mg {\\sc ii}~ core-to-wing ratio for the time interval 1980-2009. The Mg {\\sc ii}~ index shown in this figure has been derived using the NOAA Mg {\\sc ii}~ daily index (green points) version 9.1 for the period 1980 to 2002, and the SORCE satellite data (red points) for the period 2003-present. Both these time series are available from the National Geophysical Data Center website at www.ngdc.noaa.gov/stp/SOLAR/ftpsolaruv.html. Although error bars are not shown in this plot for reason of clarity, the separation between the CaK index and the corresponding Mg {\\sc ii}~ points is always less than the cumulative 1-$\\sigma$ errors. } \\label{mgii} \\end{center} \\end{figure}  Finally, visual inspection of Figure \\ref{xcak} seems to suggest a subtle but systematic decreasing in the value of the CaK index, around minima of solar activity, from 1945 to present. It is well established that chromospheric indices such as CaK index, Mg {\\sc ii}~ core-to-wing ratio, and Ly-$\\alpha$ are highly correlated with the UV and EUV irradiance. The Mg {\\sc ii}~ index, in particular, has been recently investigated to determine whether the spectral irradiance below about 300 nm shows some secular changes. \\inlinecite{2009A&A...501L..27F} has found evidence of a long-term trend in TSI, but not in solar UV irradiance. The study of the long-term properties of the UV irradiance was however limited to the investigation of the temporal behavior of the Mg {\\sc ii}~ index over the last three solar cycles, and confirmed by observations of Ly-$\\alpha$ and \\caii K over the same time period. In this paper we have the opportunity to significantly extend this time interval, by using our CaK index time series as a proxy for UV and EUV irradiance. This is validated by the high correlation between our CaK index and the Mg {\\sc ii}~ core-to-wing ratio shown in Figure \\ref{mgii}. Assuming that this correlation is preserved also for the previous cycles, shown in Figure \\ref{xcak}, we must conclude that the \"quiet\" magnetic Sun has slightly and systematically reduced its UV irradiance over the past 70 years. A regression analysis using the CaK index values from Table 1 corresponding to years of minimum in solar activity was performed to validate this conclusion. When only the last three or four solar minimum values are included in this analysis, the significance (according to a Student's t-distribution) of the calculated slope is around or well below the 95\\% confidence level, for the four and three minima respectively. This result leads to the conclusion that no significant changes have occurred in the quiet Sun UV irradiance over the last 4 decades, as also shown by other studies ({\\it e.g.} \\opencite{2009A&A...501L..27F} and references therein). However, when the last six or seven solar minimum value are considered, a significant slope- at a confidence level above 99.9\\%- is found. The value of this slope is about 0.9/year, indicating a variation of $\\approx$20-25\\% in the value of the \"quiet\" CaK index between solar cycles 17/18 and 23/24. We are not able to calculate at present how much of this variation, if any, would produce changes in the quiet-Sun's output of UV. If such a variation can be demonstrated, it could significantly impact the study of the possible effect of the Sun on the Earth's climate, as discussed in \\inlinecite{2006Natur.443..161F}. We have repeated this analysis using the smoothed temporal behavior of the daily CaK index instead of the annual means, and reached the same conclusions. Further investigation and comparison with other data sets are required, however, to positively confirm this result.  \\begin{acks} This work has been supported by NSF grant to UCLA NSF ATM-0236682. \\end{acks}"
    },
    "1005/1005.1017_arXiv.txt": {
        "abstract": "The transiting exoplanet WASP-18b was discovered in 2008 by the Wide Angle Search for Planets (WASP) project.  The \\textit{Spitzer} Exoplanet Target of Opportunity Program observed secondary eclipses of WASP-18b using \\textit{Spitzer}'s Infrared Array Camera (IRAC) in the 3.6 {\\micron} and 5.8 {\\micron} bands on 2008 December 20, and in the 4.5 {\\micron} and 8.0 {\\micron} bands on 2008 December 24.  We report eclipse depths of \\math{0.30\\pm{0.02}\\%, 0.39\\pm{0.02}\\%, 0.37\\pm{0.03}\\%, 0.41\\pm{0.02}\\%}, and brightness temperatures of 3100\\pm{90}, 3310\\pm{130}, 3080\\pm{140} and 3120\\pm{110} K in order of increasing wavelength. WASP-18b is one of the hottest planets yet discovered - as hot as an M-class star.  The planet's pressure-temperature profile most likely features a thermal inversion. The observations also require WASP-18b to have near-zero albedo and almost no redistribution of energy from the day-side to the night side of the planet.  \\comment{\\if\\submitms y \\else \\hfill\\herenote{DRAFT of {\\today} \\now}. \\fi } ",
        "introduction": "\\label{intro}   \\if\\submitms y \\clearpage \\fi \\newcommand\\s{@{\\hspace{5pt}}} \\atabon\\begin{\\widedeltab}{l\\s c\\s c\\s c|l\\s c\\s c|l\\s c\\s c|l\\s c\\s c} \\if\\submitms y \\tabletypesize{\\tiny} \\fi \\tablecaption{\\label{tab:run} Run Parameters} \\tablewidth{0pt} \\tablehead{ \\multicolumn{2}{c}{} & \\multicolumn{2}{c}{\\math{\\lambda}}   & \\multicolumn{3}{c}{Pre-Observation}  & \\multicolumn{3}{c}{Main Observation} & \\multicolumn{3}{c}{Post-Observation} \\\\ \\colhead{Obs.} & \\colhead{Date} & \\multicolumn{2}{c}{({\\microns})} & \\colhead{Start\\tablenotemark{a}} & \\colhead{Frame Time (s)} & \\colhead{Frames} & \\colhead{Start\\tablenotemark{a}} & \\colhead{Frame Time (s)} & \\colhead{Frames} & \\colhead{Start\\tablenotemark{a}} & \\colhead{Frame Time (s)} & \\colhead{Frames} } \\startdata CH13    & 2008-12-20      & 3.6 & 5.8     & 820.57796     & 2, 12  & 260    & 820.61869     &   2, 12     & 1148  & 820.79721  & 2, 12   & 10   \\\\ CH24    & 2008-12-24      & 4.5 & 8.0     & 824.35796     & 12     & 185    & 824.38431     &   2, 12     & 1148  & 824.56278  & 2, 12   & 10 \\enddata \\tablenotetext{a}{BJD - 2,454,000 ephemeris time} \\end{\\widedeltab}\\ataboff \\if\\submitms y \\clearpage \\fi \\placetable{tab:run}   Among the more than 500 extrasolar planets discovered to date\\footnote[1]{For an up-to-date listing, see http://www.exoplanet.eu}, the over 100 close-orbiting gas giants that transit their host stars have provided the most valuable clues to their physical natures. The geometry of the transit gives a direct measurement of the density of the host star and the surface gravity of the planet \\citep{SeagerEtal2003apjPlanetStarParameters,SouthworthEtal2007mnrasHD209458b}. These measurements can be combined with an estimate of the star's mass and radius to provide estimates of the planet's mass, radius and density. The closest-orbiting planets attain dayside temperatures high enough to give observable secondary eclipses at thermal-infrared wavelengths as they pass behind their host stars (e.g., \\citealp{CharbonneauEtal2005apjTrES1,DemingEtal2005natHD209458b,SingEtal2009aaGroundbased}).  Observations with the \\textit{Spitzer Space Telescope} have revealed that transiting gas giant planets can be divided into two classes based on their infrared spectral energy distributions.  A subset of the very hottest planets, with dayside temperatures in excess of 2000 K, display molecular features of CO and H\\sb{2}O in emission rather than absorption, indicating the presence of a temperature increase with height in the planet's photospheric layers.  Such temperature inversions have been inferred from the flux ratios between the four bandpasses of the Infrared Array Camera (IRAC, \\citealp{FazioEtal2004apjsIRAC}) aboard \\textit{Spitzer}, centered at 3.6, 4.5, 5.8, and 8.0 {\\micron} (channels 1--4), in the hot planets HD209458b \\citep{BurrowsEtal2008apjHD209458b, KnutsonEtal2008apjHD209458b}, XO-1b \\citep{MachalekEtal2008XO-1b}, TrES-2b \\citep{ODonovanEtal2010apjTres-2b}, TrES-4b \\citep{KnutsonEtal2009apjTres4}, XO-2b \\citep{MachalekEtal2009XO-2b} and WASP-1b \\citep{WheatleyEtal2010}, among others.  This phenomenon has been attributed to the presence of strongly-absorbing species such as TiO and VO remaining in the gas phase in the upper atmosphere, leading to the formation of a stratospheric temperature inversion \\citep{BurrowsEtal2008apjHD209458b, FortneyEtal2008TiOVO}. This led \\citeauthor{FortneyEtal2008TiOVO} to propose a scheme in which planets are assigned to class pM (with hot stratospheres) or pL (without) according to temperature, by analogy with the stellar M and L spectral classes.  \\citet{KnutsonEtal2010apjStellarActivity} suggest instead that some planets lack inversions because the TiO and VO are destroyed by UV radiation, while \\citet{SpiegelEtal2009apjTiO} claimed that TiO/VO may not be sufficiently abundant in the upper atmospheres to produce the required inversions.  \\citet{ZahnleEtal2009apjphotochemistry} proposed additional absorbers, like sulphur species, which could potentially form thermal inversions.  WASP-18b is a prime candidate for secondary-eclipse observations to test for the presence of a hot stratosphere \\citep{HellierEtal2009natWASP18b}. It orbits an F6-type star with a period of just 0.94 days and is expected to attain a blackbody equilibrium temperature approaching 2400 K, assuming zero albedo and efficient transport of heat from the dayside to the nightside of the planet. With less efficient transport, an even higher dayside temperature is expected. The planet is unusual because of its very high mass, \\math{M\\sb{\\rm p}=10.43\\pm{0.30}} Jupiter masses, \\math{M\\sb{\\rm Jup}}, and modest radius, \\math{R\\sb{\\rm p}=1.165\\pm{0.055}} Jupiter radii, \\math{R\\sb {\\rm Jup}} \\citep{SouthworthEtal2009apjWASP18b}, which give it a surface gravity an order of magnitude greater than that of any known transiting planet likely to belong to the pM class. The pressure scale height in its photospheric layers should therefore be an order of magnitude smaller than for planets of similar temperature, providing an important new dimension for tests of planetary atmospheric models in the presence of strong external irradiation. The small but precisely known orbital eccentricity \\math{e}=0.0085 \\math{\\pm0.0008} \\citep{Triaudetal2010apjSpinOrbitMeas} may impart slightly faster-than-synchronous rotation, which could help to redistribute heat from the dayside to the nightside of the planet.  WASP-18b was therefore selected as a candidate for observation as part of our \\textit{Spitzer} Exoplanet Target of Opportunity Program shortly after its discovery was confirmed. In Sections 2 and 3 of this paper we describe the observations and the analysis of the data. In Section 4 we discuss the constraints imposed on the thermal structure of the planetary atmosphere by the IRAC fluxes, in Section 5 we compare the eccentricity and orientation of the orbit from the timing and duration of the secondary eclipse with the values derived from radial-velocity observations, and in Section 6 we present our conclusions. ",
        "conclusions": "\\label{sec:concl}  \\textit{Spitzer} observed two secondary eclipses of WASP-18b using all four channels of the IRAC instrument. A blackbody model fits the observed brightness temperatures relatively well.  Slightly better fits to both inversion and non-inversion models exist with the inversion model somewhat preferred.  Because the planet is so much brighter than its predicted equilibrium temperature for uniform redistribution, the model requires near-zero albedo and very low day-night energy redistribution. The very small scale height makes this atmosphere interesting as an extreme example among irradiated planets. The addition of secondary eclipse data also improves the orbital parameters, confirming a slight eccentricity.  Files containing the lightcurves, model fits, source centers, and other ancillary data appear as electronic attachments to this article."
    },
    "1005/1005.3823_arXiv.txt": {
        "abstract": "{We present preliminary results about the detection of high redshift (U)LIRGs in the Bullet cluster field by the PACS and SPIRE instruments within the \\herschel\\, Lensing Survey (HLS) Program.  We describe in detail a photometric procedure designed to recover robust fluxes and deblend faint \\herschel\\, sources near the confusion noise. The method is based on the use of the positions of \\spitzer/MIPS 24\\mic\\, sources as priors.  Our catalogs are able to reliably (5$\\sigma$) recover galaxies with fluxes above 6 and 10~mJy in the PACS 100 and 160\\mic\\, channels, respectively, and 12 to 18 mJy in the SPIRE bands.  We also obtain spectral energy distributions covering the optical through the far-infrared/millimeter spectral ranges of all the \\herschel\\, detected sources, and analyze them to obtain independent estimations of the photometric redshift based on either stellar population or dust emission models. We exemplify the potential of the combined use of \\spitzer\\, position priors plus independent optical and IR photometric redshifts to robustly assign optical/NIR counterparts to the sources detected by \\herschel\\, and other (sub-)mm instruments.} ",
        "introduction": "Based on the IRAS, ISO, and \\spitzer\\, missions, we know that luminous infrared galaxies (LIRGs) experienced significant evolution in the past 8~Gyr, roughly decreasing their typical infrared luminosity [L(IR)] by an order of magnitude from z$\\sim$1 to z$=$0 \\citep[][]{1999ApJ...517..148F,2001ApJ...556..562C, 2005ApJ...630...82P,2005ApJ...632..169L,2009arXiv0910.5649R, 2009A&A...496...57M}.  By z$\\sim$1, they dominated the star formation activity of the Universe, being responsible for more than 50\\% of the cosmic SFR density up to at least z$\\sim$3, and playing an important role in the formation of massive galaxies at z$\\gtrsim$1.5, which exhibit high star formation efficiencies \\citep{2008ApJ...675..234P,2008ApJS..175...48R, 2009ApJ...705..617D}, strong dust obscured nuclear activity \\citep{2007ApJ...670..173D,2008ApJ...687..111D}, or both. Beyond z$\\sim$1, sub-mm surveys have also detected a population of LIRGs and (mostly) ULIRGs, identified as very active stages in the formation of massive galaxies \\citep{1997ApJ...490L...5S, 1998Natur.394..241H, 2003Natur.422..695C, 2005ApJ...622..772C,2006MNRAS.370.1185P}.  Although very useful for measuring the obscuration of optical light by interstellar dust (very abundant in the most extreme starbursts), the analysis of IR data to understand the formation and evolution of galaxies in the Universe has been classically hampered mainly by two issues.  First, the sensitivity of IR telescopes has increased significantly from one mission to the next (and also their angular resolution, although at a lower rate), but there are still problems correctly identifying all sources detected in IR surveys, most noticeably at longer wavelengths where confusion remains severe (e.g., at 70/160\\mic\\, in the \\spitzer/MIPS surveys and in the SCUBA or LABOCA sub-mm data).  Moreover, finding counterparts and getting reliable redshifts from spectroscopy for high redshift (U)LIRGs is challenging since a good fraction of them are extremely faint at optical/NIR wavelengths \\citep{2002MNRAS.337....1I, 2008MNRAS.390.1061W,2004ApJS..154..137F}. A second problem for the study of 0$<$z$\\lesssim$4 (U)LIRGs is the availability of only one photometric data point (MIPS 24\\mic) for most of them.  The fits to dust emission models to calculate L(IR), and from this derive a SFR (\\citealt{1998ARA&A..36..189K}), need a considerable extrapolation, and thus the SFR is subject to significant systematic uncertainties \\citep[][]{2006ApJ...640...92P, 2007ApJ...670..173D,barro10}.  In addition, the limited FIR data already available for galaxies at z$\\gtrsim$2 indicate that they have different dust properties from those observed for nearby (U)LIRGs \\citep{2005ApJ...622..772C,2007ApJ...668...45P, 2008ApJ...675..262R,2009ApJ...692..556R}.      The ESA \\herschel\\, Space Observatory \\citep{pilbratt10} alleviates these two issues. Thanks to its sensitivity and angular resolution at 100$<$$\\lambda$$<$500\\mic, it enables us to robustly characterize high-z galaxies. In addition to the populations of IR galaxies newly discovered by \\herschel, it will supply up to six photometric data points in the IR range for 0$<$z$\\lesssim$4 galaxies detected by MIPS at 24\\mic.  These data can be used to constrain the fits to the dust emission templates.  Even more novel and relevant is the combination of \\herschel\\, observations with the MIPS data (and also IRAC) to diminish the effects of source confusion in IR surveys. The new \\herschel\\, data indeed allow an easier and more reliable identification of counterparts from one band to the adjacent one as we move to shorter wavelengths. In addition, the \\herschel\\, bands effectively fill the gap between MIPS and (sub-)mm surveys (with SCUBA2, LABOCA, or AzTEC). By combining all these datasets, we can now use the several IR photometric data points to estimate a photometric redshift based on dust emission alone, and compare it with the estimations based on UV-to-NIR data. This can be regarded as an extension of the radio-to-IR photo-z method \\citep{1999ApJ...513L..13C,2000MNRAS.319..813D,2001PASJ...53..433R, 2003MNRAS.338..733B,2003MNRAS.342..759A,2009ApJ...694.1517D}, benefiting from a finer sampling of the emission from dust at different temperatures. The procedure can be used to improve the cross-correlation with optical/NIR galaxy samples, or even to obtain the only redshift measurement possible for very faint (high-z) IR-bright optically-faint sources. These sources may be numerous in the \\herschel\\, surveys, because of its impressive sensitivity at wavelengths probing the peak of the dust emission, which can be brighter than the optical emission by orders of magnitude.   In this Letter, we perform a preliminary analysis of the \\herschel\\, Lensing Survey (HLS, PI Egami) data taken for the Bullet cluster to show the utility of position priors and photometric redshifts obtained from UV-to-NIR and IR-to-mm data to improve the identification of \\herschel\\, and sub-mm sources with those detected at shorter wavelengths. A correct identification of IR sources will be extremely relevant to all \\herschel\\, cosmological surveys, but is of critical importance for clusters, because of their crowded nature and the possible presence of high-z lensed and distorted sources with very faint optical counterparts that may even only be detectable in the FIR/mm range. ",
        "conclusions": ""
    },
    "1005/1005.3538_arXiv.txt": {
        "abstract": "The central dominant galaxies in galaxy clusters constitute the most massive and luminous galaxies in the Universe. Despite this, the formation of these brightest cluster galaxies (BCGs) and the impact of this on the surrounding cluster environment remain poorly understood. Here we present multi-wavelength observations of the nearby poor X-ray cluster MZ\\,10451, in which both processes can be studied in unprecedented detail. {\\em Chandra} observations of the intracluster medium (ICM) in the cluster core, which harbors two optically bright early-type galaxies in the process of merging, show that the system has retained a cool core and a central metal excess. This suggests that any merger--induced ICM heating and mixing remain modest at this stage. Tidally stripped stars seen around either galaxy likely represent an emerging intracluster light component, and the central ICM abundance enhancement may have a prominent contribution from {\\em in situ} enrichment provided by these stars. The smaller of the merging galaxies shows evidence for having retained a hot gas halo, along with tentative evidence for some obscured star formation, suggesting that not all BCG major mergers at low redshift are completely dissipationless. Both galaxies are slightly offset from the peak of the ICM emission, with all three lying on an axis that roughly coincides with the large-scale elongation of the ICM.  Our data are consistent with a picture in which central BCGs are built up by mergers close to the cluster core, by galaxies infalling on radial orbits aligned with the cosmological filaments feeding the cluster. ",
        "introduction": "The most massive galaxies in the Universe occur in rich groups and clusters. These brightest cluster galaxies (hereafter BCGs) often reside close to the peak of the diffuse X-ray emission from the intracluster medium (ICM; \\citealt{jone84,zabl98,lin04,raff08}), and can have recessional velocities indistinguishable from the cluster mean (\\citealt{quin82,zabl90}; but see also \\citealt{cozi09}). Their major axis is preferentially aligned with that of the host cluster, as defined by the projected distribution of cluster galaxies (e.g., \\citealt{bing82}) or the ICM X-ray emission \\citep{hash08}, an effect which seems independent of cluster richness \\citep{full99}. Furthermore, there is a clear relationship between the stellar luminosity of central BCGs and the total mass and X-ray luminosity of their host cluster \\citep{edge91,brou08,hans09,mitt09}.  These observations suggest that many BCGs reside near the centers of the cluster gravitational potential, and that the formation of central BCGs is intimately linked to the formation and evolution of the cluster itself, as anticipated in a hierarchical structure formation scenario. This idea is supported by results of cosmological simulations which indicate that central BCGs are built up through a series of mergers near the cluster core (e.g, \\citealt{dubi98,conr07,delu07}). Such simulations also suggest that these mergers must have been predominantly dissipationless, at least for BCG major mergers occurring at low redshift (e.g., \\citealt{khoc03,delu07}). The notion that BCGs have a unique formation history is further reinforced by findings indicating that their properties are separate from those of other (cluster) galaxies. For example, the optical luminosity of BCGs is not drawn from the same distribution as that of other cluster members \\citep{trem77,vale08}, BCGs tend to have spatially extended stellar halos not seen around other galaxies (e.g., \\citealt{gonz03}), and they are more likely to be radio-loud than other galaxies of similar stellar mass \\citep{best07}.  Unfortunately, an observational understanding of the formation of BCGs is largely limited by the fact that few systems in the process of forming have actually been identified. Multiple nuclei are common in BCGs at low and moderate redshifts, suggesting that galaxy--galaxy mergers must be frequent (e.g., \\citealt{laue88, mulc06, jelt07}). In a few cases, multiple galaxies at the initial stages of merging have been identified at the centers of X-ray clusters \\citep{yama02,nipo03,tran08}, including one case in which tidal features provide direct evidence for galaxy--galaxy interactions \\citep{rine07}. However, as these examples all occur at moderate to high redshifts, detailed studies of the BCG formation process have not been possible. Potential examples at lower redshift include known dumbbell galaxies \\citep{greg92}, some of which constitute the brightest galaxy in X-ray clusters. However, at least some of these are likely to represent various stages in the mergers of {\\em cluster} cores, each already containing a fully formed central dominant galaxy \\citep{trem90,pimb08}. As such, they may not offer a particularly clean view of the build-up of the BCG itself.  For these reasons, it also remains uncertain to what extent BCG formation may have affected the thermodynamic history of baryons in X-ray bright cluster cores. For example, major galaxy--galaxy mergers near the core of a cluster could provide a mechanism for distributing metals into the ICM \\citep{zari04} and so provide a partial explanation for the pronounced central metal excesses observed in many systems (e.g., \\citealt{fino00,degr04,lecc08}). Such mergers might also act to disrupt the formation of cool cores, and thereby help explain the puzzling dichotomy between ``cool-core'' and ``non cool-core'' groups and clusters (e.g., \\citealt{sand09}).  Here we present a {\\em Chandra} study of a recently discovered low-redshift ($z\\simeq 0.06$) X-ray cluster, MZ\\,10451, with a BCG in the process of forming. To the best of our knowledge, and if excluding dumbbell galaxies but including major merger candidates in low-redshift SDSS groups and clusters \\citep{mcin08}, this represents the first known case of an ongoing central BCG major merger in an X-ray cluster in the nearby Universe. The goal of the present investigation is two-fold. We aim to establish how the ongoing galaxy--galaxy interaction is affecting the thermodynamics and chemical properties of the ICM in the cluster core. The relatively low ICM temperature of this system, $T \\approx 1$~keV, makes it attractive for a study of this kind, because any X-ray signatures of the interaction are likely to be more pronounced than in much more massive clusters. Secondly, we also hope to shed light on the BCG formation process itself, e.g.\\ by exploring whether the merging galaxies contain any hot (or cold) gas (i.e.\\ whether the merger is dissipationless), and whether they show evidence for strong nuclear activity induced by the interaction.  We assume $H_0=73$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m=0.27$, and $\\Omega_\\Lambda=0.73$. The target redshift of $z=0.0607$ then corresponds to a luminosity distance of $D \\approx 260$~Mpc, and $1'$ on the sky to 67~kpc. We further adopt Solar abundances from \\cite{ande89}, and a local Galactic H{\\sc i} absorbing column of $N_{\\rm H}=1.7\\times 10^{20}$~cm$^{-2}$ \\citep{dick90}. Unless otherwise stated, uncertainties are given at the 68\\% confidence level. ",
        "conclusions": "The $z\\approx 0.06$ low-mass cluster MZ\\,10451 is, to our knowledge, a unique system. It is X-ray bright, harbors a central BCG currently being formed from the major merger of two optically luminous early-type galaxies, and resides at a distance where the nature of the BCG merger and its impact on the surrounding ICM can be studied in unprecedented detail. We have combined {\\em Chandra} observations of the ICM in the cluster core with {\\em XMM} data of the large-scale ICM properties, and with optical, infrared, and UV data of the merging galaxies, to explore the impact of the interaction on the ICM and on the interacting galaxies themselves.  The data indicate detectable ICM emission out to at least $r \\approx 6'$, centered close to the two interacting galaxies. The ICM surface brightness in MZ\\,10451 is symmetrically distributed on both small and large scales, consistent with the expectation for a virialized system. Hence, the ongoing central galaxy--galaxy interaction has not significantly disturbed the ICM density distribution even in the group core. On the largest scales, the ICM distribution is elongated in the north-south direction, however.  The ICM exhibits a cool core and a central metal excess (here extending to $r\\sim 30$~kpc), as typically seen in undisturbed X-ray bright systems with a central BCG.  This implies either that cool cores and central abundance peaks are not necessarily destroyed by central galaxy--galaxy mergers in clusters, or that any significant central ICM heating and mixing takes place at post-coalescence stages.  The brighter of the merging galaxies presents a pronounced plume of tidally stripped stars. The ICM metallicity in the region covered by this feature is significantly enhanced. This suggests efficient local enrichment by intracluster stars ejected by the interaction, or merger-induced stripping of highly enriched gas from either galaxy. Such processes may thus have contributed to building up the pronounced central abundance peak seen in this and other systems. Some of the tidally stripped stars may represent the emergence of an intracluster light component, suggesting that such a component can form in systems that do not already host a central dominant galaxy.  Both the merging galaxies are spatially offset from the ICM X-ray peak, but only by $\\sim 10$ and $\\sim 30$~kpc, and their radial velocities are both consistent with the group mean within the errors. This demonstrates that central BCG formation can occur close to the center of the cluster potential, consistent with the notion that central BCGs are built up through a series of mergers in cluster cores.  The merger axis extends toward the X-ray peak and roughly coincides with the axis defined by the large-scale elongation of the ICM. Combined with the above, this provides support for a picture in which BCG mergers are caused by galaxies falling along radial orbits aligned with the accreting filaments that feed cluster growth.   A rough estimate based on their current separation suggest that the interacting galaxies are still $\\sim 0.5$~Gyr from nuclear coalescence. Despite being involved in a galaxy--galaxy interaction in the core of an X-ray bright system, at least one of the merging galaxies shows clear evidence of having retained an extended hot gas halo at this stage. The inferred infrared, UV, and X-ray properties of the galaxy pair are consistent with the presence of low-level obscured star formation in at least one of them. This would be consistent with a picture in which early-type galaxy interactions initially trigger residual star formation, followed by an AGN-dominated phase closer to coalescence. The tentative evidence for ongoing star formation in one of the merging galaxies suggests that low-redshift BCG major mergers need not be completely dissipationless. However, while the X-ray data reveal no evidence for strong moderately obscured AGN activity in either galaxy, the possibility of an AGN hidden behind a Compton-thick screen cannot be ruled out as an alternative to the presence of cold gas. Deeper radio observations would be required to decisively settle this issue.   Our forthcoming {\\em HST} observations of MZ\\,10451 should further elucidate the optical properties of the central merging galaxies, and thus help constrain the properties of the merger product itself. These data will also help in establishing the age of tidally stripped stars, and the generation and properties of any intracluster light in this unusual system. A detailed analysis of these and our Magellan optical data will be presented in a forthcoming paper."
    },
    "1005/1005.3154_arXiv.txt": {
        "abstract": "{That the rotation curves of spiral galaxies are generally flat in the outer regions is commonly considered an indication that galaxy disks are embedded in quasi-isothermal halos. In practice, disk-halo decompositions of galaxy rotation curves are performed in a parametric way by modeling the halo force contribution by means of expressions describing approximately the properties of the regular isothermal \\textit{sphere} or other \\textit{spherical} density distributions suggested by cosmological simulations.} {We construct self-consistent models of \\textit{nonspherical} isothermal halos embedding a zero-thickness disk, by assuming that the halo distribution function is a Maxwellian. The method developed here can be used to study other physically-based choices for the halo distribution function and the case of a disk accompanied by a bulge.} {The construction is performed by means of an iterative procedure that generalizes a method introduced in the past. In a preliminary investigation we note the existence of a fine tuning between the scalelengths $R_{\\Omega}$ and $h$, respectively characterizing the rise of the rotation curve and the luminosity profile of the disk, which surprisingly applies to both high surface brightness and low surface brightness galaxies. This empirical correlation identifies a much stronger conspiracy than the one required by the smoothness and flatness of the rotation curve (\\textit{disk-halo conspiracy}).} {The self-consistent models are found to be characterized by smooth and flat rotation curves for very different disk-to-halo mass ratios, hence suggesting that conspiracy is not as dramatic as often imagined. For a typical rotation curve, with asymptotically flat rotation curve at $V_{\\infty}$ (the precise value of which can also be treated as a free parameter), and a typical density profile of the disk, self-consistent models are characterized by two dimensionless parameters, which correspond to the dimensional scales (the disk mass-to-light ratio $M/L$ and the halo central density) of standard disk-halo decompositions. We show that if the rotation curve is decomposed by means of our self-consistent models, the disk-halo degeneracy is removed and typical rotation curves are fitted by models that are below the maximum-disk prescription. Similar results are obtained from a study of NGC 3198. Finally, we quantify the flattening of the spheroidal halo, which is significant, especially on the scale of the visible disk (\"abridged\").} {} ",
        "introduction": "\\label{Sect:Introduction}  The study of the rotation curve of spiral galaxies (from the case of M31, van de Hulst et al. 1957, to the systematic investigation of nearby galaxies in the 70's) has led to convincing evidence for the presence of dark matter halos (e.g., see Roberts 1976). To this discovery, HI rotation curves, which often sample regions well beyond the bright optical disk, played a decisive role (van Albada et al. 1985; van Albada \\& Sancisi 1986).  In this context, it has been suggested that the existence of a halo of dark matter introduces a problem of fine tuning or ``conspiracy\" (van Albada \\& Sancisi 1986). The coexisting dark and visible components, while expected to have different formation and evolution histories (even if naturally coupled by their mutual gravitational interaction), must always cooperate to produce a gravitational field inversely proportional to the radius, so that the resulting rotation curve remains flat from the inner regions (generally thought to be dominated by the visible matter) out to large radii (where the field is dominated by the diffuse dark halo). This puzzling feature is accompanied by a serious problem of {\\it degeneracy}. In fact, from a detailed study of the so-called disk-halo decomposition, it has been shown that, in general, the available photometric profile and the corresponding velocity data-points can be modeled in many different ways, either by heavy or by light disks, through an ill-determined mass-to-light ratio $M/L$ for the disk. In general, the available data are only able to set an upper limit to such $M/L$, identifying the so-called maximum-disk solution, for the model that maximizes the contribution of visible matter. Typically, the maximum-disk solution is able to fit the rotation curve out to the optical radius (Kalnajs 1983; Kent 1986).  It was immediately realized that the flatness of the outer rotation curves suggests that the dark halo embedding the visible disk might have a structure similar to that of the classical isothermal sphere, because outside the central core such sphere is characterized by a density distribution declining approximately as $\\rho \\sim 1/r^2$. Curiously, the ubiquitous presence of gravitational fields consistent with a flat rotation curve also came from studies of elliptical galaxies, not only in the nearby universe (for NGC 4472, see lower left frame of Fig.~7 of Saglia et al. 1992 and Bertin 1993; for IC 2006, see Franx et al. 1994; for a more recent systematic study, see Gerhard et al. 2001), but also for ellipticals at cosmologically significant distances (e.g., see Barnab\\'e et al. 2009, and references therein).  Of course, deviations from the paradigmatic case of galaxies such as NGC 3198 are not uncommon (e.g., see Casertano \\& van Gorkom 1991). But it has been pointed out that such deviations are generally associated with well-defined features present in the distribution of visible matter or the presence of non-circular motions (Sancisi 2004). On the other hand, pure disks characterized by the absence of specific photometric features are quite common (see also Kautsch 2009).  The literature on the general topic of the distribution of dark matter in spiral galaxies is enormous and deals with a variety of important aspects of the problem. In recent years these studies have generally focused on the requirements of current cosmological scenarios. In this Introduction it would be impossible to do justice to all the relevant studies in this general research area. However, it appears that one simple aspect of the problem has gone practically unnoticed (but see some discussion given by van Albada \\& Sancisi 1986). This is the issue of the mutual interactions between halo and visible matter in determining the self-consistent equilibrium of the basic state. In their review, van Albada \\& Sancisi (1986) addressed this point, with the help of N-body simulations, but did not carry out a systematic investigation.  Here we wish to reconsider this problem, starting from an extremely simple physical (and somewhat academic) description of the dark halo. We will thus address the problem of determining the nonspherical structure of a dark halo, assumed to be perfectly isothermal, in the presence of a given disk. We will show that even if very simple, and thus lacking some important features present in real objects (for example, the presence and role of gas are not considered explicitly in this paper, although they could be easily incorporated), this study leads to some interesting and unexpected results.  The search for equilibrium solutions of the Vlasov-Poisson equations, where the mean potential is generated by the distribution of stars populating the system, with the possible contribution of an ``external\" density distribution (e.g., associated with a co-existing dark matter halo), is a classical problem of stellar dynamics, with important analogues in the study of collisionless plasmas. In this paper we follow the approach based on the Jeans theorem to deal with the halo matter, while the disk of visible matter is taken to provide an external field. In this approach, the starting point is the identification of an appropriate form for the one-particle distribution function in phase space; in our case this is the isothermal assumption for the dark halo. The adopted distribution function then generates a $\\Phi$-dependent density in physical space, thus leading, inserted as a source in the Poisson equation, to a nonlinear problem in $\\Phi$. Here $\\Phi$ represents the gravitational potential. In the past, this method has been used frequently and has led to important insights into the dynamics of galaxies and other stellar systems. Yet the study of self-consistent equilibria in the presence of external fields that break the natural symmetry associated with the one-component problem is only rarely considered (see Bertin \\& Varri 2008). Therefore, besides the natural application to the discussion of the structure of dark halos around galaxies, this paper has some interest of its own, as a study in stellar dynamics.  The structure of the paper is the following. In Sect.~2 we comment on the so-called disk-halo conspiracy. Here we identify an empirical correlation which brings out a new and stronger conspiracy than the one associated with the flatness of rotation curves. In Sect.~3 we formulate the self-consistent problem of constructing an isothermal spheroid in the presence of a zero-thickness disk. The actual procedure to calculate the self-consistent models is reported in Appendix A. The properties of self-consistent isothermal halos are then described in Sect.~ 4. Then in Sect.~5, by modeling first a fiducial rotation curve and later by modeling the specific case of NGC 3198, we show that self-consistency actually removes the degeneracy that characterizes standard parametric disk-halo decompositions (with additional comments given separately in Appendix B). Finally, Sect.~6 summarizes the main results of the paper. This article is all based on the mathematically simple case of zero-thickness disks. A follow-up paper will address the corresponding study of disks with finite thickness. ",
        "conclusions": "\\label{Sect:Discussion}  This paper has addressed the construction of self-consistent models for a dark matter distribution assumed to be isothermal, in the presence of a zero-thickness disk. The models have then been applied to interpret the properties of observed rotation curves, with special attention to the issues of conspiracy and degeneracy that are known to affect their disk-halo decomposition.  We argue that what has been learned from the study of the strictly isothermal case should be applicable to the case of dark matter distributions characterized by a quasi-isothermal inner region (as is the case for the results of collisionless violent relaxation).  It would be interesting to test the results on the dark matter distributions predicted by current cosmological models, which definitely favor a density profile asymptotic to $r^{-3}$. This test could be performed with an approach similar to the one adopted in this paper, provided that we had a physically justified distribution function to start from. In particular, it would be interesting to test to what extent the cosmologically predicted distributions could be reconciled with the existence of cases, such as that of NGC 3198, in which the rotation curve remains very flat in a wide radial range.  Note that this paper has considered the impact of a \\textit{given} visible matter distribution on the distribution of dark matter. A study of the dynamical role of the dark halo in shaping the distribution of the visible matter would require considering a formation scenario that we do not wish to address here.  The main results obtained in this paper are the following:   \\begin{itemize}  \\item{As a preliminary empirical study, we have shown that \\textit{for both high surface brightness and low surface brightness disk galaxies} the scale length $h$ of the light profile of pure disks correlates with the scale length $R_{\\Omega}$ (defined in Eq.~(\\ref{romega})) that characterizes the central growth of the rotation curve. This suggests that the distribution of dark and visible matter in pure disk galaxies is homologous, i.e. characterized by similar global relations in different galaxies, even if significantly different in size and type (HSB vs. LSB galaxies). Based on the results obtained in the rest of the paper, this point is interpreted as an indication of a gravitational interplay between luminous and dark components in the innermost parts of the system. We argue that this homology is more relevant and surprising as a conspiracy than the actual flatness of the observed rotation curves, because, especially in the context of parametric models, rotation curves turn out to be almost flat for very different values of the disk-to-halo mass ratio.}  \\item{We have constructed self-consistent models of isothermal halos embedding zero-thickness stellar disks, thus obtaining fully self-consistent axisymmetric isothermal dark matter halos. These make basically a two-parameter family of models; the two parameters measure the weight of the disk and the central density of the halo (see Eqs.~(\\ref{alpha}, \\ref{beta})). A third parameter, the asymptotic value of the rotation curve at large radii, is in principle free but can often be treated as a fixed parameter, given the fact that it is generally well constrained by the observations; note that a discussion with the third parameter kept free is given in our study of NGC 3198 (see fifth item in this subsection). The adopted solution procedure can easily incorporate the presence of a zero-thickness disk of gas (with desired surface density). In addition, a similar procedure could be followed to construct models based on different choices of the dark matter distribution function. In turn, we can also argue that the qualitative conclusions obtained in this paper by calculating the properties of strictly isothermal halos should also apply to more general systems where the halo distribution, in the region of interaction with the disk, is quasi-isothermal.}  \\item{We have shown and quantified that, for astrophysically realistic values of the two free parameters of the models, the dark matter halo is significantly flattened, especially in the regions where it coexists with the bright optical disk. Inside one exponential scale length, the dark matter distribution is that of a thick disk; nevertheless, its contribution to the rotation curve is approximately linear in the same region, denoting the presence of a natural core.}  \\item{The rotation curves of the self-consistent models have been shown to be generally smooth and flat in their intermediate regions. In particular, self-consistency has been proved to provide a physical ``mechanism\" able to flatten rotation curves in the region of the visible disk. This \\textit{flattening effect} is stronger in the region of parameter space that appears to represent realistic systems. This result goes against the naive expectation that self-consistency should tend to enhance features present in the potential well and then in the related rotation curve.}  \\item{We have shown that the application of self-consistent models to decompose an observed rotation curve breaks the disk-halo degeneracy. This result has been achieved under some simplifying assumptions. First, for simplicity we assumed that the density profile of the disk is a perfect exponential (although our self-consistent models can be constructed for any zero-thickness axysimmetric density profile for the stellar disk). Second the dark matter distribution function has been assumed to be a perfect isothermal; this point can be released in favor of other physically justified choices. The disk-halo degeneracy breaking has been obtained first by studying an idealized rotation curve representative of a typical rotation curve of a pure disk spiral galaxy and then by a direct application to the classical case of NGC 3198. We have also provided arguments that support the idea that, to some extent, such degeneracy problem has been artificially created by the use of parametric models (such as the model associated with Eq.~(\\ref{vdm})). In fact, the rotation curves generated by these parametric models have been shown to be exceedingly uniform and featureless.}  \\item{We have thus derived the best-fit self-consistent disk-halo decompositions and proved that they favor mass decompositions characterized by stellar disks of intermediate weight, which are still important in the central regions of the system, but are significantly lighter than maximum-disk solutions. On the other hand, the homology property noted in the first item of these Conclusions would still make such sub-maximal solutions consistent with the small scatter of the Tully-Fisher (1977) relation. Of course, a full demonstration of this conclusion would require that we moved beyond the case of the fiducial rotation curve (see Sect.~5.1),  by studying, in the way described in Sect.~5.2 for NGC 3198, a complete sample of spiral galaxies. But this would obviously go well beyond the objectives of the present paper.}  \\item{The resulting rotation curve has been interpreted as resulting from a \\textit{suppression effect}, characteristic of self-consistency. Self-consistency would affect a proposed maximum-disk decomposition by significantly reducing the dark halo contribution to the total rotation curve and thus leading to worse fits, especially at intermediate radii. }  \\item{The resulting $M/L$ values for the stellar disk should be studied for each individual galaxy, in view of their implications and consistency with respect to what is known about the relevant stellar populations. In this context, an important test would be the study of LSB galaxies. However, this very important issue goes well beyond the goals of the present paper.}  \\end{itemize}  In a follow-up paper we will generalize the results of this paper to the case in which the disk is characterized by finite thickness, with results that may be of interest to projects such as the Disk Mass Project (Verheijen et al. 2004 and Westfall et al. 2008)."
    },
    "1005/1005.4068_arXiv.txt": {
        "abstract": "We present the WFC3 Infrared Spectroscopic Parallel (WISP) Survey. WISP is obtaining slitless, near-infrared grism spectroscopy of $\\sim 90$\\ independent, high-latitude fields by observing in the pure parallel mode with Wide Field Camera-3 on the {\\it Hubble Space Telescope} for a total of $\\sim 250$ orbits. Spectra are obtained with the G102 ($\\lambda=0.8-1.17~\\mu$m, R$\\sim210$) and G141 grisms ($\\lambda=1.11-1.67~\\mu$m, R $\\sim$130), together with direct imaging in the J- and H-bands (F110W and F140W, respectively). In the present paper, we present the first results from 19 WISP fields, covering approximately 63 arcmin$^{2}$. For typical exposure times ($\\sim6400$~s in G102 and $\\sim2700$~s in G141), we reach $5\\sigma$\\ detection limits for emission lines of $f\\sim5\\times 10^{-17}$~ergs s$^{-1}$~cm$^{-2}$\\ for compact objects. Typical direct imaging 5$\\sigma$-limits are 26.8 and 25.0 magnitudes (AB) in F110W and F140W, respectively. Restricting ourselves to the lines measured with highest confidence, we present a list of 328 emission lines, in 229 objects, in a redshift range $0.3 < z < 3$. The single-line emitters are likely to be a mix of H$\\alpha$ and [OIII]5007,4959 \\AA, with H$\\alpha$ predominating. The overall surface density of high-confidence emission-line objects in our sample is approximately 4 per arcmin$^{2}$. These first fields show high equivalent width sources, AGN, and post starburst galaxies. The median observed star formation rate of our \\ha\\ selected sample is 4 M$_{\\odot}$ year$^{-1}$. At intermediate redshifts, we detect emission lines in galaxies as faint as $H\\textunderscore140 \\sim 25$, or $M_R < -19$, and are sensitive to star formation rates down to less than 1 M$_{\\odot}$ year$^{-1}$. The slitless grisms on WFC3 provide a unique opportunity to study the spectral properties of galaxies much fainter than $L^*$ at the peak of the galaxy assembly epoch. ",
        "introduction": "\\label{sec:intro}  While there is considerable uncertainty regarding the detailed star formation history of galaxies of various masses and types, it has become clear that the span of cosmic time $0.5 \\le z \\le 2.5$ encompasses most of the star formation in the history of the Universe \\citep[e.g.][]{hopkins04,daddi07}. This broad epoch is also when many galaxies were ``assembled,'' when major mergers and the accretion of satellite galaxies and gas combined to produce the variety of morphologies and structures of modern galaxies.  There is mounting evidence for an evolution of star-forming galaxies during this time period, from high-mass galaxies dominating at $z \\sim 3$, to low-mass galaxies at the present.  Deep, wide IR surveys are required to cover much of the peak epoch of star formation, because of the redshifting of the key diagnostic spectral features of stellar evolution, and to minimize the effect of obscuration by dust.  The difficulty of obtaining observed-frame near-infrared spectra causes much of our empirical perspective on galaxy evolution to be biased by the way in which spectroscopic samples are selected.  Color selection reliably selects common objects, but misses unexpected ones; magnitude-limited samples necessarily favor bright-continuum objects.  Ground-based searches for emission lines from faint high-redshift galaxies are severely impacted by the bright NIR airglow.  Excluding OH terrestrial airglow frequencies restricts redshift coverage and compromises redshift identification. Space-based spectroscopic surveys offer significant advantages for unbiased target selection. The grisms on the Near Infrared Camera and Multi Object Spectrometer (NICMOS), the Space Telescope Imager and Spectrograph (STIS) and the Advanced Camera for Survey (ACS) have demonstrated the power of spectroscopy from above the atmosphere, where the low background emission permits slitless observing \\citep{gardner98, pmc99, teplitz03a, pirzkal04,drozdovsky05}. Furthermore, the terrestrial wavelength gaps can only be filled from space, as first demonstrated by the pure-parallels program with the NICMOS G141 grism \\citep{pmc99, yan99, shim09}.  Much as ground-based objective prism surveys (e.g. Gallego et al. 1999) were used to define flux-limited samples of H$\\alpha$ emitters at $z \\sim 0$, the NICMOS grism observations produced complete samples of emission-line objects at intermediate redshifts. Yan et al. (1999) and Hopkins et al. (2000) used observations with the NICMOS H-band (G141) grism to derive H$\\alpha$ luminosity functions, and to show that the luminosity density, and hence star formation rate density, has evolved by a factor of $\\sim 10$ from $z \\sim 1.4$ to the present. Shim et al. (2009) used a larger sample to derive improved H$\\alpha$ luminosity functions in two redshift bins and confirmed the trends seen in the analysis by Yan et al. and Hopkins et al. The complete reductions/extractions of grism data for the principal modes of NICMOS were released in the Hubble Legacy Archive \\citep[HLA,][]{freudling08}. Nearly all of the published results from the NICMOS grisms were based on data with the G141 H-band disperser. The poor sentivity of the NICMOS detector at short wavelengths limited the utility of the J\\textunderscore110-band grism.  The grisms on ACS have been used to discover faint Ly$\\alpha$ emission sources and faint Lyman continuum break sources at $z \\sim 3 - 4$\\ \\citep{malhotra05,pirzkal07} as well as to derive spectrophotometric redshifts for a large sample of faint galaxies \\citep[e.g.][]{straughn08,straughn09}. Extremely deep ACS grism observations have shown the power of low resolution spectroscopy above the atmosphere to detect spectral breaks indicative of passively evolving systems \\citep[e.g.][]{daddi05}\\ or active star forming galaxies \\citep{pirzkal04,rhoads05,malhotra05,pasquali06}. The full release of the reduced ACS grism data in the HLA is underway \\citep{kuemmel09b}. The STIS slitless modes work at higher resolution than the NICMOS and ACS grisms, but the relatively low sensitivity of STIS has limited its ability to sample populations beyond the reach of ground-base spectroscopic surveys \\citep{teplitz03b}.  Wide Field Camera Three (WFC3) on Hubble has a larger field of view, significantly better detectors and superior sampling compared to camera 3 on NICMOS. The net improvement in survey speed is approximately a factor of 20.  This enhanced survey speed does not fully capture the gains provided by WFC3. The quality of the spectra, in terms of resolution, a well-focused point spread function, and uniform sensitivity are also greatly superior to those offered by NICMOS.  In HST's Cycle 17, we began a new observing program of WFC3 pure parallels, the WFC3 Infrared Spectroscopic Parallel (WISP) survey.  In $\\sim 250$\\ orbits, the program will obtain near-infrared slitless spectroscopy of $\\sim 90$\\ uncorrelated high latitude fields (PI=Malkan, GO-11696). This will produce a large sample of emission-line galaxies to measure the star formation rate continuously from at $z \\sim 0.5$ to $z \\sim 2.5$, over which ground-based searches are severely limited. The study of dust and metallicity in a large, unbiased sample of galaxies and the evolution of the star formation density are among the main goals of our project. In this paper, we present the analysis of the first 19 fields, from which we have selected high-confidence emission lines.  In Sections \\ref{sec:obs}\\ and \\ref{sec:datared}, we describe the survey and data reduction.  We discuss the results in Section \\ref{sec:results}, and finally offer some thoughts on the future possibilities in Section \\ref{sec:discuss}.  Throughout, we assume a $\\Lambda$-dominated flat universe, with $H_0=71$\\ km s$^{-1}$\\ Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.73,$\\ and $\\Omega_{m}=0.27$. All magnitudes are in AB system. ",
        "conclusions": "\\label{sec:discuss}  We have presented the first analysis of the WISP Survey of slitless spectroscopy with the NIR channel of WFC3, taken in pure parallel mode.  In the first 19 fields, we surveyed 63 arcmin$^{2}$. At the highest confidence level we have detected 229 emission line objects.  We reach typical $5\\sigma$\\ sensitivities of $f\\sim5\\times 10^{-17}$~ergs s$^{-1}$~cm$^{-2}$\\ for compact lines.   While wide-field narrow-band imaging surveys from the ground \\citep[e.g.][]{sorbal09,momcheva10} yield samples of H$\\alpha$ emitters that are larger than ours and near-IR slit spectroscopy of selected targets can yield emission-line ratios to faint levels (e.g. Pettini et al. 2001), the WFC3 grisms offer a unique opportunity to obtain continuous spectral coverage in the $0.8 - 1.7$\\mic\\ region with high sensitivity. Our early results show that we can identify star forming galaxies with space densities comparable to, or larger than, those identified with color selection techniques (e.g. BM/BX or LBG selection; Adelberger et al. 2004) to similar star formation rate thresholds (Shapley et al. 2005, Erb et al. 2006). The typical emission-line galaxy at $z > 1.3$ in our sample has strong [OIII]5007,4959 emission and a large [\\ion{O}{3}]/H$\\beta$ ratio, similar to the ratios seen in the small number of LBGs at $z>1.5$ with rest-frame optical spectra (Pettini et al. 2001, Teplitz et al. 2000, Lemoine-Busserolle et al. 2003, Maiolino et al. 2008, Hainline et al. 2009). Coupled with the high [\\ion{O}{3}]/H$\\alpha$ ratio and weak [\\ion{O}{2}]3727 these suggest low metal abundances. Future papers from our survey will examine line ratios and their implications for abundances and ionization sources. With the support of our ongoing ground-based, follow-up observing programs, we will be able to measure the mass-metallicity relation at crucial intermediate redshifts, to derive the \\ha\\ luminosity function at different redshifts and to study the evolution of the star formation rate density.   Many of our fields contain emission-line objects with extremely high equivalent widths (cf. Figure \\ref{fig:examples4}). These strongly star forming emission-line dominated objects are reminiscent of the ultra-strong emission-line (USEL) galaxies at $0<z<1$\\ discovered by \\cite{kakazu07} \\citep[see also][]{hu09} and the ``green pea'' galaxies found in SDSS by Cardamone et al. (2009). These objects have metallicities comparable to the lowest observed in local sources and may be dwarf galaxies undergoing their first major star formation episode.   The sensitivity and resolution of the G102 grism also provide the opportunity to search for Ly$\\alpha$\\ emission lines in addition to the rest-frame optical emission lines discussed here. We have excluded such analysis from the present paper given the limitations of the initial, visual selection of faint emission lines. After further analysis, WFC3 will offer a significant new approach to high-redshift ($z> 6$) searches for Ly$\\alpha$. Though extensive ground-based narrowband searches for Ly$\\alpha$-emitters have been performed up to $z=7$, with the most distant confirmed source at $z=6.96$ \\citep[][]{iye06}, it is difficult beyond this redshift to find wavelengths with both high sky transmission and low sky emission. Thus, many surveys are limited to a very small range at $z=7.7$ or $z=8.8$\\ \\citep[e.g.][]{sorbal09}, without any reliable Ly$\\alpha$ emission-line candidates.  In addition to the WISP survey, two large observing programs with the WFC3 grisms will contribute significant results. First, a G141 survey of the GOODS North fields is underway (GO-11600; PI=Weiner), reaching depths comparable or slightly better than WISP.  Secondly, the CANDELS multi-cycle treasury program (GO-12060 through 12064; PI=Faber and Ferguson) will obtain deep spectroscopy of selected pointings targetted to follow-up supernovae found in direct imaging. The WISP survey has some limitations compared to these programs due to the parallel mode, including the lack of dithering, but our first analysis demonstrates the great efficiency of this use of the telescope.  % WFC3 parallels can obtain large spectroscopic samples of emission-line objects during the peak epoch of star formation, free from the most common selection biases."
    },
    "1005/1005.1960_arXiv.txt": {
        "abstract": "{} {Our aim is to understand the evolutionary sequence of high-mass star formation from the earliest evolutionary stage of high-mass starless cores, via high-mass cores with embedded low- to intermediate-mass objects, to finally high-mass protostellar objects.} {{{\\it Herschel}} far-infrared PACS and SPIRE observations are combined with existing data at longer and shorter wavelengths to characterize the spectral and physical evolution of massive star-forming regions.} {The new {{\\it Herschel}} images spectacularly show the evolution of the youngest and cold high-mass star-forming regions from mid-infrared shadows on the Wien-side of the spectral energy distribution (SED), via structures almost lost in the background emission around 100\\,$\\mu$m, to strong emission sources at the Rayleigh-Jeans tail. Fits of the SEDs for four exemplary regions covering evolutionary stages from high-mass starless cores to high-mass protostellar objects reveal that the youngest regions can be fitted by single-component black-bodies with temperatures on the order of 17\\,K. More evolved regions show mid-infrared excess emission from an additional warmer component, which however barely contributes to the total luminosities for the youngest regions. Exceptionally low values of the ratio between bolometric and submm luminosity additionally support the youth of the infrared-dark sources.} {The {{\\it Herschel}} observations reveal the spectral and physical properties of young high-mass star-forming regions in detail.  The data clearly outline the evolutionary sequence in the images and SEDs. Future work on larger samples as well as incorporating full radiative transfer calculations will characterize the physical nature at the onset of massive star formation in even more depth.} ",
        "introduction": "\\label{intro}  Characterizing the early evolutionary stages of (high-mass) star formation is difficult because the cold cores ($\\sim$10 to 20\\,K) have the peak of the spectral energy distribution (SED) at far-infrared wavelengths which were hardly accessible until recently.  {\\it Herschel} has changed this situation completely, and we are now able to study the young (massive) star-forming regions in detail. For this early {\\it Herschel} study we selected the complex associated with the High-Mass Protostellar Object (HMPO) IRAS\\,18223-1243 and the Infrared Dark Cloud IRDC\\,18223-3 at a distance of $\\sim$3.7\\,kpc \\citep{sridha}. This region hosts most evolutionary stages in massive star formation, starting with high-mass starless cores, continuing with high-mass cores with embedded low- to intermediate-mass protostars potentially forming massive stars, to HMPOs with already embedded and likely still accreting massive protostars (Fig.~\\ref{continuum}).  Previous studies focused either on the HMPO IRAS\\,18223-1243 (e.g., \\citealt{sridha,beuther2002a}) or at a high-mass core with embedded low- to intermediate-mass protostar IRDC\\,18223-3 \\citep{garay2004,beuther2007a,fallscheer2009}.  However, the whole complex with the youngest starless cores has not been studied so far. Here we combine {\\it Herschel} PACS and SPIRE data at far-infrared wavelengths \\citep{A&ASpecialIssue-HERSCHEL,A&ASpecialIssue-PACS,A&ASpecialIssue-SPIRE} with Spitzer mid-infrared data and (sub)mm continuum observations at the Rayleigh-Jeans tail of the spectrum.  \\begin{figure*}[ht] \\includegraphics[angle=-90,width=18.4cm]{14532fg1.eps} \\caption{The color-scale presents the wavelengths as marked in each panel. 70, 100 and 160\\,$\\mu$m are from PACS, 250\\,$\\mu$m is from SPIRE. The contours show the SCUBA 850\\,$\\mu$m continuum emission from 0.05 to 0.35\\,Jy\\,beam$^{-1}$ (in 0.1\\,Jy\\,beam$^{-1}$ steps) and additionally two stronger contours at 0.65 and 1.25\\,Jy\\,beam$^{-1}$.  The right panel labels which part of the region are dominated by what evolutionary stage. The circles and labels in the 160\\,$\\mu$m panel mark the sources for which SEDs were extracted (Fig.~\\ref{fig_seds}).} \\label{continuum} \\end{figure*} ",
        "conclusions": "\\label{general}  These regions represent a potential evolutionary sequence where IRDC\\,18223-south1 resembles the youngest detectable stage with a large emission peak of cold dust continuum emission on the Rayleigh-Jeans tail of the SED longward of 160\\,$\\mu$m. However, at 100\\,$\\mu$m the source is lost in the background, and at shorter wavelengths it remains an absorption shadow. The source IRDC\\,18223-south2 appears a bit more evolved because one detects an emission source from 70\\,$\\mu$m to longer wavelengths. Nevertheless, the region is still young and cold without any emission at 24\\,$\\mu$m. Continuing on the evolutionary ladder, IRDC\\,18223-3 exhibits 24\\,$\\mu$m emission but remains dark in the IRAC bands. \\citet{beuther2007a} called this a high-mass core with an embedded low- to intermediate-mass protostar potentially becoming massive at the end of the evolution.  Finally, at the northern end of the filament IRAS\\,18223-1243 comprises all features of a typical HMPO \\citep{sridha}.  This sequence is also reflected in the SED fitting. For the cold sources, the single black-body fits result in low temperatures between 16 and 18\\,K, luminosities of only a few hundred L$_{\\odot}$ and gas masses of several hundred M$_{\\odot}$.  The luminosities of these cold sources are not related to internal heating, but are mainly produced by the external radiation field.  The low temperatures are consistent with previous estimates toward infrared dark clouds based on spectral line observations (e.g., \\citealt{sridharan2005,pillai2006}). We note that our fitted temperatures are also in the range of temperatures derived from COBE for the general ISM (e.g., \\citealt{reach1995}). On the given scales ($>$0.65\\,pc) at these early evolutionary stages, the gas and dust clumps still have temperatures comparable to those of the general ISM. For the most evolved HMPO, the fitted cold temperature is warmer than for the other sources.  Furthermore, the luminosity already indicates the presence of an HMPO, while the gas reservoir is still large, allowing for further accretion.  While the total luminosity compares well with older fits that were only based on IRAS data \\citep{sridha}, the cold temperature of the new fit is lower than those only based on IRAS data (28 versus 50\\,K, respectively). This is because IRAS alone did not sample the peak of the SED well and hence overestimated the temperature. Therefore sampling the full SED is important even for such evolved sources to derive reasonable temperature estimates.  In the framework of the $L/L_{\\rm{sm}}$ ratio, the HMPO is above 100 where 200 was defined as the border between class 0 and class I for low-mass sources \\citep{andre1993}. While this cannot easily be translated to high-mass star-forming regions, it shows that even the HMPO is still in a young evolutionary stage (see \\citealt{sridha}). The three infrared-dark sources all have exceptionally low $L/L_{\\rm{sm}}$ ratios around 10, lower than the high-mass cores reported by \\citet{mueller2002}, and among the lowest values so far reported in the literature (e.g., \\citealt{young2003}). This supports the extreme youth of these sources. That we do not see a marked difference in $L/L_{\\rm{sm}}$ for the three infrared-dark sources indicates that -- in the evolution of the massive cores -- first signatures for star formation do not come from a global warm-up of the bulk of dust and gas. More evolved sources like the HMPO 18223-1243 in turn have multiple T components which show up at mid-infrared wavelengths, but most of the cold dust can still be reasonably well fitted by a higher but single temperature.  While we present here one of the first sources observed with PACS and SPIRE, there is tremendous potential in this kind of observations. For the future we anticipate full radiative transfer modeling of such regions, which will result e.g.~in more detailed temperature maps and density structures over the entire field. Furthermore, the {\\it Herschel} key project EPOS (PI O.~Krause) contains more than 40 regions in young evolutionary stages. A combined analysis of the whole sample should result in a robust characterization of the physical properties and early evolution of high-mass star formation."
    },
    "1005/1005.3376_arXiv.txt": {
        "abstract": "The determination of Li and proton-capture element abundances in globular cluster (GC) giants allows us to constrain several key questions on the multiple population scenarios in GCs, from formation and early evolution, to pollution and dilution mechanisms. In this Letter we present our results on Li abundances for a large sample of giants in the intermediate-metallicity GC NGC~6121 (M4), for which  Na and O have been already determined by Marino et al. The stars analyzed are both below and above the red giant branch bump luminosity. We found that the first and second generation stars share the same Li content, suggesting that a Li production must have occurred. This is a strong observational evidence providing support for the scenario in which asymptotic giant branch stars are GC polluters. ",
        "introduction": "\\label{sec:intro} The presence of multiple populations as a characterizing property of globular clusters (GCs) is widely accepted nowadays. Since the first photometric and spectroscopic studies, which revealed multiple sequences (e.g. Dickens \\& Woolley, 1967; Lee et al. 1999, 2009; Bedin et al. 2004) and/or light element (anti)correlations (see Gratton et al. 2004), it became evident that GC stars are neither coeval nor chemically homogeneous. Hence, GCs host at least two stellar generations. Thanks to the advent of 8$-$10 m class telescopes, which also allowed targeting fainter main-sequence (MS) stars, several abundance studies have shown that chemical (anti)correlations are also present in unevolved (MS/turnoff (TO)) or scarcely evolved (SGB) stars (Gratton et al. 2001; Carretta et al. 2004; Pasquini et al. 2005). This evidence implies that a previous generation of stars has activated CNO, NeNa, and MgAl cycles in their interiors, in order to deplete O and Mg and enhance Na and Al, respectively. The origin/nature of such kind of stars is still debated with two coexistent scenarios: (1) intermediate-mass asymptotic giant branch (AGB) stars in hot bottom burning phase (Ventura \\& D'Antona 2009) and (2) fast rotating massive stars (FRMS; Decressin et al. 2007).   Our group has recently carried out an extensive survey, focusing on the determination of proton-capture elements in 19 GCs (Carretta et al. 2009c), with the main objective of discovering and understanding the nature (and the extent) of the chemical (anti)correlations and their link with global cluster parameters (horizontal branch (HB) morphology, metallicity, etc.). The large sample of stars ($\\sim$1200 GC members), analyzed in a very homogeneous and accurate way, revealed that while the Na$-$O anticorrelation is present in all the GCs (i.e. the second generation is not a ``perturbation\"), the shape of the Na$-$O distribution varies from cluster to cluster (Carretta et al. 2009c). On the other hand, the analysis of UVES  spectra for $\\sim$ 200 stars by Carretta et al. (2009b) has shown that the Mg$-$Al anticorrelation is not present in all GCs. Both these indications suggest that the typical polluter masses change from cluster to cluster: this variation is apparently driven by a combination of cluster luminosity and metallicity.   In this context, lithium abundances offer a complementary approach to $p$-capture elements allowing to address several important issues. If no Li is produced by the polluters, the multiple population scenario predicts that Li and O are positively correlated, while Li and Na anticorrelated. Na-poor, O/Li-rich stars are the first population born in the cluster (they share the same chemical composition of field stars at the same metallicity), and Na-rich, O/Li-poor stars constitute the second generation. Within the same hypothesis, Li is an excellent tracer of the dilution process acting within each star: only through Li abundance determinations we can determine the amount of pristine (and of polluted) material present in each star. In particular, we hope to answer two fundamental questions: do 100\\% polluted stars (Li$\\sim$0) exist or even the most extreme population still contains a certain fraction of primordial matter? Also, is the minimum measurable Li content the same for all GCs or does it vary from cluster to cluster?   On the other hand, Li offers the exciting chance to observationally constrain the nature of the polluters. If the progenitors of second generation stars are FRMS, they have destroyed their original Li content. On the other hand, if AGB stars are responsible for intra-cluster pollution, they may have non-negligible Li yield, given the Li production via the ``$^7$Be transport\" mechanism (Cameron \\& Fowler 1971). As a consequence, we can reveal if AGB stars are responsible for GC pollution through two main observational evidences: (1) the presence of very Li-rich stars among GC populations and (2) the lack of Li$-$Na anti-correlation (or Li$-$O correlation), with the second generation stars also showing a rather high Li content. In a recent work, we focused on $\\sim$90 TO stars belonging to the metal-rich GC 47 Tuc: in this case, likely because of the high metallicity, a large star-to-star scatter in Li abundances erases any Li$-$Na anticorrelation, while Li and O appear to be only weakly positively correlated (D'Orazi et al. 2010). The cluster seems to display a different behaviour from NGC~6752 and NGC~6397: for the first one, Pasquini et al. (2005) detected a significant Li$-$Na anticorrelation (and also Li$-$O correlation and Li$-$N anticorrelation). Concerning the second cluster, the situation seems more complex since Lind et al. (2009) detected a quite constant value in Li abundances, with only three stars (out of 100) driving a hint of Li$-$Na anticorrelation. Enlarging the sample of simultaneous determinations of Li and $p$-capture elements in GCs is hence of paramount importance: in this Letter we present Li results on the intermediate-metallicity ([Fe/H]=-1.18; Carretta et al.\\ 2009a) cluster NGC~6121 (M4), by analyzing the same sample of a hundred red giant branch (RGB) stars studied by Marino et al.\\ (2008). As found by Marino and coworkers, M4 hosts two distinct populations of stars, mainly characterized by a different sodium content (i.e., the Na-rich and Na-poor groups), and defining different sequences on the color$-$magnitude diagram $U$ versus ($U-B$). We derived Li abundances for the stars belonging to the two groups, and located both below and above the RGB bump luminosity; this evolutionary stage plays a fundamental role in this context.  Theoretical models (Iben 1967), confirmed by observations of field stars by Gratton et al. (2000), predict a depletion in Li due to the first dredge-up (1DUP) of a factor of $\\sim$ 20 at the base of the SGB branch. On the lower RGB, below the bump luminosity, the molecular weight gradient associated with the H abundance jump acts as a barrier that prevents further extra mixing (Charbonnel et al. 1994); then the Li abundance remains constant until the RGB bump is reached. At this stage, the H shell reaches and cancels this discontinuity and non-standard mixing processes (non-convective extramixing, see, e.g., Charbonnel \\& Zahn 2007) cause the total destruction of the remaining Li. ",
        "conclusions": "\\label{sec:sum} We report in this Letter Li abundances for a sample of $\\sim$100 giants, both below and above the RGB bump luminosity, belonging to the GC NGC~6121 (M4). The main purpose of our work was the study of the correlation (if any) between Li content and elements involved in $p$-capture reactions, the last ones already presented for the very same stars by Marino et al. (2008). The principal results we obtained in our investigations can be summarized as follows.  \\begin{itemize} \\item[1.] As expected, Li tends to disappear as the stars reach the RGB bump luminosity; however the Na-rich and Na-poor stars show very different trend of Li with magnitude. Specifically, while the decline of Na-poor stars with $M_{\\rm V}$ is rather smooth, there is a very brusque decrease for Na-rich stars. The so-called ``thermohaline mixing\", which seems responsible for extramixing processes at the RGB bump, is not predicted to have different outcomes in Na-poor and Na-rich stars. Further investigations of this aspect are mandatory both from observational and theoretical points of view. New observations, focusing on Li determination along the RGB in several GCs, are necessary to assess if M4 is a ``peculiar\" case or other GCs share the same behavior; as a consequence model predictions could be revised in this sense. \\item[2.] M4 does not show any Li$-$Na anticorrelation, with first and second generation stars having almost the same Li content. Along with similarities in Li abundances, the larger scatter found in Na-rich stars, indicates that a Li production, from the previous generation of polluters, must have happened. This provides support to intermediate-mass AGB stars responsible for intra-cluster contamination, since FRMS can only destroy Li.  \\end{itemize}  \\begin{table*} \\caption{Stellar Parameters and Abundances for Our Sample Stars. This table is available in its entirety in a machine-readable form in the online journal. A portion is shown here for guidance regarding its form and content.}\\label{t:tab1} \\begin{center} \\begin{tabular}{lccccccc} \\hline\\hline star  & T$_{\\rm eff}$ & log$g$ & $\\xi$       & [Na/Fe] & err$_{\\rm Na}$ & log~n(Li) & err$_{\\rm Li}$ \\\\ &   (K)         &        &(kms$^{-1}$) &         &                &           &               \\\\ \\hline 30345 & 4850          & 2.73   &1.31         &  0.43   & 0.07 \t        &  1.37     & 0.08\\\\ 30452 & 4830          & 2.56   &1.25         &  0.06   & 0.03 \t        &  1.29     & 0.09\\\\ 30719 & 4810          & 2.65   &1.24         &  0.42   & 0.06 \t        &  1.22     & 0.08\\\\ 31306 & 4900          & 2.87   &1.33         &  0.40   & 0.05 \t        &  1.35     & 0.07\\\\ \\hline\\hline \\end{tabular} \\end{center} \\end{table*}"
    },
    "1005/1005.4410_arXiv.txt": {
        "abstract": "We study perturbations in the multi-field axion N-flation model, taking account of the full cosine potential.  We find significant differences with previous analyses which made a quadratic approximation to the potential. The tensor-to-scalar ratio and the scalar spectral index move to lower values, which nevertheless provide an acceptable fit to observation.  Most significantly, we find that the bispectrum non-gaussianity parameter $\\fnl$ may be large, typically of order 10 for moderate values of the axion decay constant, increasing to of order 100 for decay constants slightly smaller than the Planck scale. Such a non-gaussian fraction is detectable. We argue that this property is generic in multi-field models of hilltop inflation. ",
        "introduction": "The axion N-flation model is based on a set of $\\Nf$ uncoupled fields, labelled $\\phi_i$, each with a potential \\cite{DKMW} \\begin{equation} V_i = \\Lambda_i^4 \\left ( 1- \\cos \\alpha_i \\right) \\,, \\end{equation} where $\\alpha_i = 2\\pi\\phi_i/f_i$ and $f_i$ is the $i^{\\mathrm{th}}$ axion decay constant.  In a more general model couplings may exist between the fields, but we will not consider these.  The mass of each field in vacuum satisfies $m_i = 2\\pi \\Lambda_i^2/f_i$, and the angular field variables $\\alpha_i$ lie in the range $(-\\pi,+\\pi]$. Without loss of generality we will set initial conditions with all $\\alpha_i$ positive.  If only a single field is present this model is known as natural inflation \\cite{natural}.  Calculation of the observables $n$, $r$ and $\\fnl$ makes use of the $\\delta N$ formula \\cite{deltaN}, which considers how the total number of $e$-folds of expansion $N$ is modified by field perturbations.  We define slow-roll parameters for each field as \\begin{equation} \\epsilon_i \\equiv \\frac{\\Mp^2}{2} \\left( \\frac{V_i'}{V_i} \\right)^2\\,, \\label{eq:epsi} \\end{equation} where $\\Mp \\equiv (8 \\pi G)^{-1/2}$ is the reduced Planck mass, a prime denotes the derivative of a function with respect to its argument, and no summation over $i$ is implied.  The global slow-roll parameter $\\epsilon \\equiv -\\dot{H}/H^2$ can be written as a weighted sum $\\epsilon \\simeq \\sum_i (V_i/V)^2 \\epsilon_i$, in which each field contributes according to its share of the total energy density. We must have $\\epsilon < 1$ during inflation.  We work in the horizon-crossing approximation, in which the dominant contribution to each observable is assumed to arise from fluctuations present only a few $e$-folds after horizon exit of the wavenumber under discussion.  After smoothing the universe on a superhorizon scale somewhat smaller than any scale of interest, the horizon-crossing approximation becomes valid whenever the ensemble of trajectories followed by smoothed patches of the universe approaches an attractor.  We suppose that inflation exits gracefully, with each field settling into the minimum of its potential.  The horizon-crossing formulas will then be a reasonable approximation. Using this method, and conventional definitions for each observable parameter \\cite{PDP}, we find \\begin{eqnarray} \\Ps_\\zeta & = & \\frac{H_*^2}{4\\pi^2} \\sum_i N_{,i}N_{,i} = \\frac{H_*^2}{8\\pi^2 \\Mp^2} \\sum_i \\frac{1}{\\epsilon_i^*}\\,; \\label{eq:P} \\\\ n-1 % & = &-2\\epsilon_* - \\frac{8\\pi^2}{3H_*^2} \\sum_j \\frac{\\Lambda_j^4}{f_j^2} \\frac{1}{\\epsilon_j^*} \\Big/ \\sum_i \\frac{1}{\\epsilon_i^*} \\,; \\label{eq:n} \\\\ r & = & \\frac{2}{\\pi^2 \\Ps_\\zeta} \\frac{H_*^2}{\\Mp^2} = 16 \\Big/ \\sum_i \\frac{1}{\\epsilon_i^*}\\,; \\label{eq:r} \\\\ \\frac{6}{5} \\fnl & \\simeq & \\frac{\\sum_{ij}N_{,i}N_{,j}N_{,ij}}{\\left(\\sum_k N_{,k}N_{,k}\\right)^2} = \\frac{r^2}{128} \\sum_i \\frac{1}{\\epsilon_i^*} \\frac{1}{1+\\cos \\alpha_i^*} \\,, \\label{eq:fnl} \\end{eqnarray} where $N_{,i}$ and $N_{,ij}$ are respectively the first and second derivatives of $N$ with respect to the fields, and $\\ast$ indicates evaluation at horizon crossing (determined by Eq.~(\\ref{eq:ntot}) below).  In writing Eq.~\\eqref{eq:fnl} any intrinsic non-gaussianity among the field perturbations at horizon crossing has been neglected, a good approximation provided $\\fnl > 1$ \\cite{Maldacena,VW}.  Our sign convention for $\\fnl$ matches WMAP \\cite{WMAP7}, and the non-gaussianity is predicted to be of local type. The observed amplitude of perturbations is obtained by adjusting the $\\Lambda_i$ to give an appropriate value of $H_*$.  Under a quadratic approximation to each potential, it can be shown that Eqs.~\\eqref{eq:r} and \\eqref{eq:fnl} recover their single-field values of order $\\sim 1/N_\\ast$ \\cite{single,VW}, making $\\fnl$ undetectably small.  The spectral index can be shown to be less than its single field value $1-2/N_*$ \\cite{LR} with equality only in the equal-mass case. Its value for a given choice of parameters must be computed numerically \\cite{KLs}.  However, we will see that these results all change whenever our initial conditions populate the hilltop region. ",
        "conclusions": "We have described a new mechanism for generating observably-large cosmic non-gaussianity, based on the strongly-dispersive dynamics of fields in a hilltop region. In a multi-field context such as the axion N-flation model, assisted inflation can yield a viable spectral index without a major dilution of the non-gaussianity. As compared to the quadratic potential approximation to N-flation, we found a substantial decrease in $r$, a modest increase in $1-n$, and a substantial increase in $\\fnl$.  These changes will happen whenever initial conditions have a significant probability of populating the hilltop region, such as the uniform (in field angle) initial conditions we chose.  Searches have previously been made for models which achieve $|\\fnl| \\gg 1$ while preserving slow-roll during inflation \\cite{BCH}. The N-flation model is of this type, but offers several advantages.  The non-gaussian fraction is naturally bounded above, so that $\\fnl$ cannot become arbitrarily large.  Therefore our predictions do not depend on a sudden exit from inflation, e.g.\\ triggered by a hybrid transition, to prevent $\\fnl$ from growing to an unacceptable value. Equally important, our large signal does not derive from a singularity of the $e$-folding history $N$, as a function of its initial conditions.  These means we can rely on a perturbative expansion. We can simultaneously satisfy observational constraints on the spectral index and tensor fraction.  Moreover, this result seems generic. Inflation is self-replicating on top of the hill, sometimes described as `topological inflation' \\cite{topo}.  Coupled with the dispersion of trajectories originating from the vicinity of the hilltop, this implies large non-gaussianity may not be uncommon over a landscape of scalar field vacua."
    },
    "1005/1005.3006_arXiv.txt": {
        "abstract": "We present a morphological analysis of the rest-frame ultraviolet emission of 78 resolved, high signal-to-noise $z \\sim  3.1$ Lyman Alpha Emitters (LAEs) in the Extended Chandra Deep Field South (ECDF-S).  Using {\\it HST/ACS\\/} $V$-band images taken as part of the Galaxy Evolution from Morphology and SEDS (GEMS), Great Observatories Origins Deep Survey (GOODS), and Hubble Ultra Deep Field (HUDF) surveys,  we investigate both single component and multi-component LAEs, and derive concentration indices, S\\'ersic indices, ellipticities, and half-light radii for all resolved components and systems with a signal-to-noise $S/N > 30$.  We show that, although the LAE population is heterogeneous is nature, most Ly$\\alpha$ emitters highly concentrated, with a distribution of $C$ values similar to that measured for field stars; this suggests that the diagnostic is a poor discriminator near the resolution limit.  The LAEs also display a wide range of S\\'ersic indices ($0 < n < 12$), similar to that seen for galaxies in the local neighborhood. However, the majority of LAEs have $n < 2$, and a visual inspection of the images suggests that the small$-n$ objects have extended or multimodal luminosity profiles, while the LAEs with $n > 2$ have compact components surrounded by diffuse emission. Moreover, unlike nearby spiral galaxies, whose distribution of ellipticities is flat, the LAE ellipticity distribution peaks near $1 - b/a \\sim 0.55$.   Thus, the population has more in common with $z \\sim 3$ Lyman-break galaxies than local star-forming objects. ",
        "introduction": "In the local universe, most galaxies fall onto the Hubble Sequence \\citep{HubbleSeq} which runs from older, red, quiescent galaxies with a $r^{1/4}$ surface-brightness profiles to star-forming, gas-rich disks with exponential luminosity distributions.  These light profiles are generally described by the \\citet{Sersic} parameter,  defined through $I(r) \\propto r^{1/n}$, where $n$ typically ranges from $\\sim 1$ (an exponential disk) to $\\sim 4$ (an r$^{1/4}$-law spheroid). Out to intermediate ($z \\sim 1.5$) redshifts, galaxies can generally be placed onto the local Hubble Sequence quite reliably \\citep{Cons04}.  However, at larger distances, this sequence breaks down, as most sources appear clumpy and irregular \\citep[e.g.,][]{Steidel96,Papovich05,Conselice05,Venemans05,Pirzkal07} and are therefore difficult to classify.   Consequently, although the morphological parameters of distant galaxies can be quantified via the use of codes such as {\\small GALFIT} \\citep{GALFIT} and {\\small GIM2D} \\citep{GIM2D}, these parameters do not necessarily directly translate into the familiar Hubble sequence seen in the nearby universe. Moreover, the robustness of the derived parameters is highly dependent on the signal-to-noise ($S/N$) of the measurement,  with simulations suggesting that a $S/N$ of at least 15 is required \\citep{Ravindranath06} to reliably fit a high-redshift galaxy.  To address some of these difficulties, a number of non-parametric measures of galaxy morphology have been developed, including the concentration, asymmetry, and clumpiness indices \\citep[CAS;][]{CAS}, and the Gini coefficient, which measures a profile's departure from uniformity \\citep{Lotz04}.  These systems have been used extensively on samples of high-redshift galaxies, since they require no a priori knowledge about the functional form of the luminosity distribution. In addition, unlike other fitting techniques which use smooth two-dimensional profiles, these non-parametric techniques can also quantify the non-uniformity of a galaxy (i.e., test for the presence of star-forming regions), and define its asymmetry. Their one drawback, of course, is that the values one obtains may depend on the image depth, and that this sensitivity has not been well-quantified for high-redshift galaxies.  The majority of galaxies known at high ($z \\geq 2.5$) redshifts have been identified via the Lyman-break technique, wherein galaxies with Lyman-limit discontinuities are identified by their unique location in color-color space \\citep{Steidel96}. Morphological analyses of these Lyman Break Galaxies (LBGs) in the rest-frame UV have revealed that most of these systems are disturbed and disk-like (i.e., with exponential profiles), and only $\\sim 30\\%$ have profiles consistent with galactic spheroids \\citep[e.g.,][]{Ferguson04,Lotz06,Ravindranath06}. Moreover, \\citet{Ravindranath06} have found that the distribution of ellipticities derived from GALFIT are skewed towards higher values, and that this asymmetry becomes more pronounced at higher redshift. This dominance of ``elongated'' morphologies seen in LBGs has been interpreted as evidence for galaxy formation within filamentary structure or the presence of proto-bars that span the entire visible disk of the galaxy.  Ly$\\alpha$ Emitters(LAEs) are a class of high-redshift galaxies which have not been studied as extensively as LBGs.   Like their Lyman break counterparts, LAEs between $2 < z < 4$  are actively star-forming objects \\citep[e.g.,][]{CowieHu}, and at the bright end, the sample of LAEs overlaps that of LBGs.   However, a typical Ly$\\alpha$ emitting galaxy has  a lower stellar mass ($\\sim 10^9 M_{\\sun}$),  higher mass-specific star-formation rate ($\\sim 10^{-8} M_{\\sun}$~yr$^{-1} M_{\\sun}^{-1}$, and lower dust content ($E(B-V) \\sim 0.1$) \\citep{Venemans05,Gawiser07,lai} than any galaxy discovered via the Lyman-break technique.  Moreover, LAEs are less clustered than LBGs, suggesting a lower dark-matter halo mass, and their bias factor $b \\sim 1.7$) consistent with that expected from the progenitors of present-day $L^*$ galaxies \\citep{Gawiser07,guaita}.   To date, relatively few LAEs have been studied morphologically.  In the rest-frame UV, most $3 \\lesssim z \\lesssim 6$ LAEs are small ($\\lesssim 1$~kpc), compact ($C>2.5$), and barely resolved at {\\it HST\\/} resolution \\citep{Venemans05,Pirzkal07,Overzier08,taniguchi}.  Those LAEs that are not compact  ($\\sim 20 - 45$\\%) exhibit clumpy or irregular morphology, with diffuse components extending to several kiloparsecs. Due to their small sizes and low luminosities, however, categorizing these objects within the context of existing classification schemes has been difficult.   Here we present a study of the rest-frame ultraviolet morphology of a statistically-complete sample of $z=3.1$ LAEs \\citep{GronwallLAE} present on archival {\\it HST}/ACS\\/ $V$-band images from the Galaxy Evolution from Morphology and SEDS \\citep[GEMS,][]{GEMS}, Great Observatories Origins Deep Survey \\citep[GOODS,][]{GOODS}, and Hubble Ultradeep Field \\citep[HUDF,][]{HUDF} surveys. In Paper I \\citep{Bond09}, we investigated the size distribution of LAEs by developing an analysis pipeline which identifies each component of an LAE, and measures its size and brightness.  We found that most LAEs are extremely compact, with half-light radii, $r_e < 2$~kpc, and that a $S/N > 30$ is required for a robust measure of morphology.  Here, we extend this study, and provide more detailed morphological information on each object.   In \\S~\\ref{sec:data} and \\ref{sec:method}, we describe the data and detail the algorithms used in our morphology-fitting pipeline analysis.   In \\S~\\ref{sec:results}, we present the best-fitting morphological parameters for each galaxy, and explore how these parameters vary with image depth.  Finally, in \\S~\\ref{sec:discussion}, we discuss the implications of our findings.  Throughout this paper, we will assume a concordance cosmology with $H_0=71$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m=0.27$, and $\\Omega_{\\lambda}=0.73$ \\citep{WMAP}.  With these values, $1\\arcsec = 7.75$~physical~kpc at $z=3.1$. ",
        "conclusions": "\\label{sec:discussion}   From Figures~\\ref{fig:HUDFPanels} and \\ref{fig:GOODSPanels} it is clear that our sample of $z\\sim3$ LAEs have % a broad range of morphologies.  While most sources are % small ($r_e < 2$ kpc) and concentrated ($2 < C < 4$), the rest-frame UV emission from a typical LAE is neither smooth nor symmetrical.   LAE ellipticities are skewed towards higher values of $\\epsilon$, in a manner similar to that seen for LBGs at the same redshift \\citep{Ravindranath06}. This suggests the presence of  elongated, prolate structures, which may be indicative of merging activity or ``chain\" star formation. Unfortunately, it is difficult to probe the concentration indices of the population: even for the subset of LAEs that {\\small GALFIT} calls resolved, the distribution of $C$ values is similar to that measured for stars.  Next generation space-based imagers are needed to properly probe this parameter.    Our LAEs also possess a broad range of S\\'ersic indices (see Figure~\\ref{fig:sersic}). As in the local universe \\citep{blanton}, the distribution of $n$ values peaks near $\\sim 1$ with a broad distribution with values ranging from $n \\lesssim 1$ and $n \\gtrsim 12$.  Our population does not, however, have the high-$n$ population associated with luminous, red galaxies in the local universe. Whether we are observing the formation of the Hubble sequence is not  clear, but there does appear to be a qualitative change in morphology at $n \\sim 2$.  At greater $n$ values, LAE components appear more compact, relative to their surrounding diffuse emission, than objects with $n \\sim 1$.  However, these ``spheroid-like'' LAEs could also be the result of compact star-forming regions imbedded in the diffuse emission of a disk-like or irregular galaxy.  Deeper imaging is required to make this distinction.   If we classify LAEs with $n > 2.5$ as ``bulge-like\" and LAEs with $2.5 \\geq n > 0.8$ as ''disk-like\" we find that 31\\% of our sample of resolved LAE components with $S/N > 30$ fall into the former category, while 44\\% belong to the latter. In a sample of LBGs at the same redshift, \\citet{Ravindranath06} found 27\\% were bulge-like and 42\\% were disk-like.   From the data, it appears that the light profiles of LBGs and LAEs at this redshift have similar distributions.  A subset of our LAEs consist of multiple components, and this can be interpreted as evidence for merger activity. When interacting systems are in close proximity, they may appear as as one single object with a small value of $n$ (see, for example, LAE $25$ in Figure~\\ref{fig:HUDFPanels}).  If we follow \\citet{Ravindranath06} and consider sources with $n < 0.8$ (i.e., with light profiles that are flatter than exponential) to be merger candidates, then $\\sim 24$\\% of our sample would be classified as such.  This again is consistent with the 31\\% value found by \\citet{Ravindranath06} for LBGs at a similar redshift. We note that in a stacked sample of $z  \\sim5.7$ LAEs created by \\citet{taniguchi}, $n \\sim 0.7$ which may indicate a larger fraction of merger candidates in LAE samples at higher redshift.   The sample of Taniguchi, however, covers a significantly brighter range of Ly$\\alpha$ luminosity ( $3 \\times 10^{43}$  to $6.3 \\times 10^{42}$ ergs~sec$^{-1}$) than our sample at $z\\sim3.1$ ($\\sim 10^{43}$ to $1.3 \\times 10^{42}$ ergs~sec$^{-1}$).   Ideally, such a comparison between morphological properties of LAE samples would be done by comparing samples at similar luminosities, but such samples do not currently exist besides that presented here.    Of course, estimating merger rate using measured $n$ values fails to account for LAEs in which the interacting galaxies are at large separation.  In these cases, SExtractor will consider the objects separate as components (see, for example, LAE $11$ in Figure~\\ref{fig:GOODSPanels}).  On the other hand, galaxies with clumpy internal star formation can be mistaken for mergers when viewed in the rest-frame UV\\null.  Clumpy star-formation will lead also to low measured values of $n$ \\citep[e.g.,][]{rawat09} and hence an overestimate of the estimated merger rate. To rule out this possibility, and to make progress towards determining % the true LAE merger fraction, we would require deep $V$-band imaging ($m_{lim,5\\sigma} \\sim 30$) for a much larger sample of LAEs than is covered by the HUDF\\null.  Our morphological results for LAEs at $z\\sim3$ are broadly similar to those found for continuum-selected LBG objects at the same redshift. This consistency supports the contention of \\citet{law} and \\citet{pentericci} that that the presence of Ly$\\alpha$ in emission has little bearing on rest-frame UV morphology of $z \\sim 3$ star-forming galaxies. Comparable results were also found by \\citet{taniguchi} at $z\\sim6$, and may support a scenario in which the observed differences between LAEs and LBGs are driven by variations in dust content and HI gas \\citep{verhamme}. On the other hand \\citet{vanzella} found that dropout galaxies at $z\\sim4$ with Ly$\\alpha$-emission were significantly more concentrated than the rest of the sample.     Disentangling the effects of emission-line strength on morphology as a function of redshift will require substantially larger samples with deep high-resolution imaging at multiple redshifts, preferably tracking the same rest-frame continuum wavelength free of contamination from the Ly$\\alpha$ emission line.  Our results for the morphologies of $z \\sim 3$ LAEs highlight the need for consistency between measurements made at different survey depths and the need for systematic checks to verify the robustness of derived parameters.  It also points out the importance of making measurements consistently over a wide range of redshifts. This paper represents the first attempt to measure morphological parameters of a large sample of LAEs as previous work has concentrated primarily on measuring sizes \\citep{Venemans05,Bond09}, used small ($n \\lesssim 12$) sample sizes \\citep{Overzier08,Pirzkal07}, or measured the profiles of a stacked sample due to low signal-to-noise \\citep{taniguchi}.    The advent of larger samples of LAEs from redshifts ranging from $z=2.1$ to $z>7$ should allow for meaningful comparisons across redshifts down to similar luminosity limits to be made. Without such data, it is impossible to form a clear picture of how the process of galaxy formation proceeds over time.    The measurements presented here are limited by the fact that we are observing the morphologies of these high-redshift galaxies in the rest-frame ultraviolet where individual clumps of on-going star-formation may strongly affect the light distributions \\citep[e.g.,][]{overzier10}.  Future work in the rest-frame optical (observed infrared) using WFC3 on {\\sl HST} should alleviate some of these concerns, and allow us the bulk of the LAE stellar population.  A comparison of the two bandpasses will then enable us to examine how star formation proceeds in these objects."
    },
    "1005/1005.1529_arXiv.txt": {
        "abstract": "We consider the evolution of primordial black holes by including nonstationarity in the formation process and accretion of radiation in Brans-Dicke theory. Specifically, we focus on how $\\eta$, the fraction of the horizon mass the black hole comprises capturing nonstationarity, affects the lifetimes of these primordial black holes. Our calculation reveals that the primordial black hole dynamics is controlled by the product $f\\eta$ where $f$ is the accretion efficiency. We also estimate the impact of $\\eta$ through  $f\\eta$ on the primordial black holes' initial mass fraction constraint obtained from the $\\gamma$-ray background limit. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.1003_arXiv.txt": {
        "abstract": "We compare the expected abundance of cosmic-ray electrons and positrons from pulsars and magnetars.  We assume that the distribution of infant pulsars and magnetars follows that of high-mass stars in the Milky Way and that the production rate of cosmic rays is proportional to the spin-down and magnetic-decay power of pulsars and magnetars, respectively.  In combination with primary and secondary cosmic-ray leptons from other sources (especially supernova remnants), we find that both magnetars and pulsars can easily account for the observed cosmic-ray spectrum, in particular the dip seen by HESS at several TeV and the increase in positron fraction found by PAMELA. ",
        "introduction": "Introduction}   Recent measurements of cosmic rays from energies of one GeV to one TeV have shown a spectrum \\citep{2008Natur.456..362C,2008PhRvL.101z1104A} and abundance of positrons \\citep{2009Natur.458..607A} that are difficult to reconcile with astrophysical sources.  These observations have naturally stirred up great interest among the particle physics community and especially people who study dark matter \\citep[e.g.][]{2009NatPh...5..176H}. Essentially the problem is that the observed spectrum is much harder than that expected to be produced by astrophysical sources such as supernova shocks or pulsars, so it is quite natural to look at more exotic possibilities, including decaying dark matter. On the other hand, several works have also examined the question of whether or not astrophysical sources can account for the observed spectrum and positron abundance by including a reasonable model for the diffusion of cosmic rays through the Galaxy. In this Letter we consider the expected abundance of cosmic-ray electrons and positrons (hereafter referred to as cosmic rays) produced by magnetars and pulsars.  Magnetars \\citep{Dunc92,Thom95} are ultramagnetized neutron stars with fields $B_* \\sim 10^{14} - 10^{16}$ G, fueled by either magnetic field decay \\citep{Thom96,Heyl98decay} or residual thermal energy \\citep{Heyl97kes,Heyl97magnetar}.  Magnetars are thought to produce copious numbers of electron-positron pairs both in bursts \\citep{Thom95,Heyl03sgr} and in quiescence \\citep{Heyl05sgr,2007ApJ...657..967B}. Moreover, the local radiation from these pairs is thought to power the bursts from soft-gamma repeaters \\citep{Thom95} as well as the recently discovered non-thermal emission from soft-gamma repeaters and anomalous x-ray pulsars in quiescence \\citep{2005A&A...433L...9M,2006ApJ...645..556K}. Consequently, magnetars seem a natural candidate to make an important contribution to the observed abundance of electrons and positrons observed at Earth.  \\citet{2008arXiv0812.4457P} presented a detailed calculation of the expected spectrum of cosmic rays in various models for their production from pulsars. His calculation focused on the population of gamma-ray pulsars. The situation for magnetars is somewhat different because these objects are much rarer than gamma-ray pulsars; the birth rate for magnetars is roughly ten times smaller than that of pulsars.  Consequently, the cosmic ray flux that we observe can come both from the magnetars that are active today as well as from more ordinary looking objects that may have been active magnetars in the past such as RXJ~0720.4-3125 \\citep{Heyl98rxj,Heyl98decay}.  To calculate the expected abundances of cosmic rays from magnetars, we generate a Monte Carlo simulation of the magnetar population over the past ten million years and compare it with a similar simulation of the pulsar population over the past million years. ",
        "conclusions": "We have derived the expected spectrum and composition of cosmic rays from magnetars under the reasonable assumptions that the decay of the external magnetic field powers the cosmic-ray production and that the birthrate of magnetars is about ten percent of that of pulsars.  We assume that pulsar spin-down powers the cosmic-ray production from these more weakly magnetized neutron stars.  We find that either the ensemble of magnetars or that of pulsars can reproduce many of the puzzling features of the spectrum of electron-positron cosmic rays recently discovered by ATIC and HESS.  The differences in the appearance and location of these features in the cosmic-ray abundance between the models results mainly from the the relative rarity of magnetars.  Although the relative heights of the two bumps in cosmic rays detected at Earth do depend on the the assumed power-law production spectrum, the location of the transition between them does not. On the other hand, the ensemble of pulsars cannot easily reproduce the sharp dip in abundance around 2~TeV found by HESS -- pulsars do reproduce the upturn at 5~TeV, so a combination of the two ensembles may explain both the dip and the upturn.  Both pulsars and magnetars can account for the positron abundance in cosmic rays with energies of about one to one-hundred GeV. With additional modelling of the cosmic-ray diffusion, energy losses and production, even more detailed agreement with the observations can be achieved."
    },
    "1005/1005.2236_arXiv.txt": {
        "abstract": "We present a new formula to numerically construct configurations in rotational equilibrium, which consist of multiple layers. Each layer rotates uniformly or differentially according to cylindrical rotation-laws that are different from layer to layer. Assuming a different barotropic equation of state (EOS) for each layer, we solve the Bernoulli equation in each layer separately and combine the solutions by imposing continuity of the pressure at each boundary of the layers. It is confirmed that a single continuous barotropic EOS is incompatible with the junction condition. Identifying appropriate variables to be solved, we construct a convergent iteration scheme. For demonstration, we obtain two-layered configurations, each layer of which rotates rapidly with either an \"$\\Omega$-constant law\" or a \"$j$-constant law\" or a \"$v$-constant  law\". Other rotation laws and/or a larger number of layers can be treated similarly. We hope that this formula will be useful in studying the stellar evolution in multi-dimension with the non-spherical configuration induced by rotation being fully taken into account. ",
        "introduction": "\\label{sec:Intro}  It is well known that stars are generically rotating on the main sequence and, in particular, massive stars are rapid rotators~\\citep{Fukuda:1982,Tassoul:2000}. Although the distribution of angular momentum in the stellar interior is poorly known except for the sun, it is expected that the inner portion is rotating more rapidly than the outer part as the star evolves and the central part of the star contracts. In fact some recent theoretical studies on the evolution of rotating stars have demonstrated that massive stars in their late evolutionary phases develop a central core that is rotating more rapidly than the outer envelopes~\\citep{Heger:2000ud,Hirschi:2004ks, Limongi:2000km}. Hence the differential rotation is supposed to exist quite commonly in the stellar interior especially at the advanced stages of evolution.  The above-mentioned works on the evolution of rotating stars ignore non-spherical deformations of stars and the angle-averaged centrifugal force is added as a correction to the spherical models. This is not justified, however, if the star is rotating rapidly and the physical conditions on the rotation axis and on the equatorial plane are substantially different. Then the rotational equilibrium should be properly taken into account. This will be particularly important for the investigation of the progenitors of gamma ray bursts, since they are supposed to be driven by the gravitational collapse of very rapidly rotating massive stars~\\citep{Woosley:2005gy}.  Over the years substantial effort has been made to numerically obtain configurations in rotational equilibrium in various contexts. Beginning with the pioneering works by \\citet{Ostriker:1968}, a robust iterating formula was developed by~\\citet{Hachisu:1986} and was extended to general relativistic and/or magnetized stars~\\citep{Bocquet:1995,Cook:1993qj,Kiuchi:2008ch,Komatsu:1989,Tomimura:2005}. The polytropic equation of state (EOS) was replaced by more generic ones~\\citep{Cook:1993qr,Kiuchi:2007pa,Kiuchi:2009zt}. One of the limitations of these studies that hamper the application to the study on the evolution of rotating stars, putting aside the treatment of convections and meridional circulations, is the assumption that the whole star is rotating cylindrically, that is, the angular velocity is constant on each of the concentric cylinders and, as a result, the EOS is barotropic, that is, the pressure is a function of density alone~\\citep{Tassoul:2000}. On the contrary, the theoretical studies based on the spherical models indicate clearly that the stellar core and envelopes composed of different elements rotate rather independently of each other, since the steep gradient of mean molecular weight tends to suppress the transport of angular momentum beyond the boundary of elements. If one attempts to employ the rotational equilibrium configurations in the study of the post-main-sequence evolutions of massive stars, therefore, it is almost mandatory to treat multiple layers that obey different rotation-laws.  Motivated by these facts, we present a new formulation to numerically construct non-relativistic configurations in rotational equilibrium, which consist of multiple layers. We assume that each layer rotates still cylindrically but the rotation-law, namely the angular velocity as a function of the distance from the rotation axis, can be different from layer to layer. This assumption allows us to employ the conventional formula based on the Bernoulli equation in each layer, a big advantage over the original partial differential equations. The EOS should be barotropic in each layer accordingly. We then introduce a junction condition, that is, the continuity of pressure at the boundary of the layers to combine them in such a way that the whole star is in rotational equilibrium. In this paper, mainly for demonstration purposes, we obtain two-layered configurations with each layer having different polytropic EOS's for three representative rotation-laws:(1) $\\Omega$-constant law (rigid rotation), (2) $j$-constant law and (3) $v$-constant law (see the next section for the exact definitions of these rotation laws). In principle, there is no problem in treating different rotation-laws and/or a larger number of layers.  The paper is organized as follows. In Sec.~\\ref{sec:basic}, we describe the new formulation to obtain rotational equilibrium configurations with multiple layers, explaining numerical issues in detail. Section~\\ref{sec:res} is devoted to the demonstration of some model computations. In Sec.~\\ref{sec:sum}, we give some discussions and summarize the paper. ",
        "conclusions": "\\label{sec:sum}  Bearing in mind the application to the study of rotational massive stars in their late evolutionary phase, in this paper we have proposed a new formula to construct multi-layered configurations in rotational equilibrium. This is an extension of the Hachisu self-consistent field scheme that is based on the Bernoulli equation and meant originally for single-layered configurations that are rotating cylindrically with a barotropic EOS. In our method, on the other hand, each layer is assumed to rotate still cylindrically with a barotropic EOS but the rotation laws and EOS's are different from layer to layer. We have shown that the pressure should be continuous at the layer boundary whereas the density is in general discontinuous across the boundary, which is an alternative demonstration that the EOS cannot be identical for the adjacent layers. We have identified the variables that are appropriate to make the iteration scheme convergent. This is indeed a crucial ingredient in our formula.  For demonstration, we have actually constructed several configurations with two layers for three representative rotation laws, which we have referred to as the $\\Omega$-constant, $j$-constant and $v$-constant laws in this paper. We have found that the virial equation is satisfied with a typical error of $10^{-5}$ irrespective of the rotation laws if we deploy $1000 \\times 200$ mesh points and we have also demonstrated that the error is reduced as the number of mesh points is increased. Incidentally, it has been confirmed that a non-rotational configuration is also reproduced by the present scheme. From these results it is obvious that our method works well and is robust indeed. The application of the present formula to more realistic problems will be published elsewhere~\\citep{Nagakura:2010}.  As commented in Sec.~\\ref{sec:basic}, it is straightforward to extend our scheme to the configurations with more than two layers though the procedure becomes a bit more involved. Although we have combined the rotation laws of the same family but with different parameters for simplicity in this paper, two rotation laws of different families can be treated in the same way. The implementation of more realistic EOSs will pose no problem in principle as long as they are barotropic. We may employ the idea by \\citet{Jackson:2005,MacGregor:2007zy} that the pressure, density, and temperature are assumed to be functions of the effective potential alone. Moreover, the present formula will be able to treat configurations with a topology of torus by relaxing the assumption that the surface extends itself to the symmetry axis and by choosing appropriate points on the equator to impose the conditions corresponding to Eqs.~(\\ref{eq:cond4}), (\\ref{eq:cond3}) and (\\ref{eq:cond6}) although we do not know how realistic such configurations are.  In our formula, the layer boundary is determined from the condition that the pressure be continuous there. In reality, however, the layers in the stellar interior correspond to the regions of different chemical compositions and their boundaries are determined by the thermodynamical conditions for nuclear burnings. This difference originates from the fact that we have imposed piece-wise cylindrical rotation laws. In the actual stellar interior, each layer obeys a boroclinic EOS and, as a result, rotates non-cylindrically. Moreover, the gas motions in the meridian section such as convections and meridional circulations are likely to exist generically. Then the original partial differential equations should be solved somehow, which is a formidable task and will need an entirely new approach. Our formula, therefore, is admittedly a rather crude approximation to the reality but, hopefully, not so bad one if one chooses an appropriate rotation law for each layer. In fact, it will be much better than any approximate configurations with only a single-layer.   The real challenge will be to somehow implement chemical evolutions to the sequence of rotational configurations. One possibility may be an extension of the idea employed in most of the current one dimensional evolution models of rotational massive stars~\\citep{Heger:2000ud,Hirschi:2004ks,Limongi:2000km}. Under the assumption that the thermodynamical conditions as well as the chemical abundances are uniform  on each surface of constant effective potential, we solve the nuclear network locally and then transfer generated energy on the multi-dimensional mesh. The transport of angular momentum may be also approximated by diffusion. Since the resultant distributions of thermodynamical quantities and elements will in general be non-uniform on the surface of constant effective potential, we will take their angular averages on the surface and solve the new rotational equilibrium for the obtained equations of state and rotation law. This completes the single cycle and the iteration of this process will give the temporal evolution of rotational stars. We hope that this procedure is feasible and that the formulation presented in this paper will contribute to the study of the influences of non-sphericity on the evolution of rapidly rotating massive stars."
    },
    "1005/1005.0233_arXiv.txt": {
        "abstract": "We study the spherical collapse model for several dark energy scenarios using the fully nonlinear differential equation for the evolution of the density contrast within homogeneous spherical overdensities derived from Newtonian hydrodynamics. While mathematically equivalent to the more common approach based on the differential equation for the radius of the perturbation, this approach has substantial conceptual as well as numerical advantages. Among the most important are that no singularities at early times appear, which avoids numerical problems in particular in applications to cosmologies with dynamical and early dark energy, and that the assumption of time-reversal symmetry can easily be dropped where it is not strictly satisfied. We use this approach to derive the two parameters characterising the spherical-collapse model, i.e.~the linear density threshold for collapse $\\delta_\\mathrm{c}$ and the virial overdensity $\\Delta_\\mathrm{V}$, for a broad variety of dark-energy models and to reconsider these parameters in cosmologies with early dark energy. We find that, independently of the model under investigation, $\\delta_\\mathrm{c}$ and $\\Delta_\\mathrm{V}$ are always very close to the values obtained for the standard $\\Lambda$CDM model, arguing that the abundance of and the mean density within non-linear structures are quite insensitive to the differences between dark-energy cosmologies. Regarding early dark energy, we thus arrive at a different conclusion than some earlier papers, including one from our group, and we explain why. ",
        "introduction": "Over the last decade, a wealth of evidence was accumulated in favour of the conclusion that the expansion of our Universe is accelerating, mainly from the observation of the type-Ia supernovae \\citep{Riess1998, Perlmutter1999, Kowalski2008} and the cosmic microwave background (CMB) \\citep{Komatsu2009} in combination with measurements of the Hubble constant and large-scale structures (LSS) \\citep{Cole2005}. Assuming the validity of general relativity on large scales, a possible explanation for this accelerating expansion, is obtained by introducing a component of the cosmic fluid, the dark energy, with equation-of-state parameter $w<-1/3$.  Despite efforts from both observational and theoretical sides, the nature of the dark energy remains obscure. Consequently, a plethora of different models has been proposed for the origin and the time evolution of the dark energy, see for example \\cite{Copeland2006} for a comprehensive review. The simplest model assumes that the dark energy is connected with the vacuum energy, the so called cosmological constant, with equation-of-state parameter $w=-1$. Despite the fact that observations constrain the present value of $w$ quite tightly, the time evolution of the equation-of-state parameter is rather poorly constrained. Thus, it is natural to study more general classes of models allowing a time evolution of the dark-energy component, such as models involving scalar fields.  Scalar fields occur naturally in particle physics and in string theory and could thus be candidates to explain the nature of the dark-energy if they are sufficiently strongly self-interacting. This class includes quintessence models, phantom models, K-essence, tachyon models and so forth. Scalar fields are described by their Lagrangian density with a kinetic term $\\dot{\\phi}^2/2$ and a potential term $V(\\phi)$. The equation-of-state parameter then follows from the canonical energy-momentum tensor. If the dark energy is spatially homogeneous, $P=\\dot{\\phi}^2/2+V(\\phi)$ and $\\rho=\\dot{\\phi}^2/2-V(\\phi)$, giving $w=P/(\\rho c^2)$.  The cosmological-constant case is recovered if the kinetic energy is negligibly small compared to the potential energy. Dark energy affects first of all the expansion rate, causing geometrical effects that can be revealed through distance measurements, such as the luminosity distance to distant supernovae. Secondly, it affects structure formation, the early stages of which can be quantified by the growth factor. Thus, structure formation will be affected by the amount of dark energy and by its dynamical evolution over cosmic history.  One model recently suggested as a candidate for solving the fine-tuning problem of the cosmological constant was the class of early dark-energy cosmologies (EDE) \\citep{Wetterich2004,Doran2006}, according to which the contribution of dark energy at early times is not negligible. Thus, to produce the same amount of structure now, structure formation should start earlier and proceed more slowly than in the common $\\Lambda$CDM model. It would thus compensate the additional, opposing effects of the early dark-energy contribution. An analytic calculation based on the evolution equation for the radius of a spherical, homogeneous perturbation and various assumptions \\citep{Bartelmann2006} implied a substantial increase in the number of objects compared to a standard $\\Lambda$CDM model. While this expectation was confirmed by \\cite{Sadeh2007}, subsequent $N$-body simulations by \\cite{Francis2009b} and \\cite{Grossi2009} found instead that the effect of EDE on the mass function of dark-matter haloes and its evolution is almost negligible: The EDE class of models predicts differences in the mass function of only a few percent with respect to the $\\Lambda$CDM model. A new numerical derivation of the linear overdensity parameter, also based on the differential equations of the spherical collapse model \\citep{Francis2009a} was in perfect agreement with the numerical simulations.  Motivated in part by this discrepancy of results derived from the same model, we are here addressing the problem of determining the time evolution of the linear overdensity $\\delta_c(z)$ in a completely different way, using a perturbative approach based on Newtonian hydrodynamics directly. One of the advantages, apart from increased numerical stability, is that no time-reversal symmetry needs to be assumed for the spherical collapse.  The structure of the paper is as follows. In Sect.~\\ref{sect:nhydro}, we present the basic equations of Newtonian hydrodynamics and sketch the derivation of the equations used to obtain the linear over density threshold $\\delta_{\\mathrm{c}}$ and the non-linear virial overdensity parameter $\\Delta_{\\mathrm{V}}$. In Sect.~\\ref{sect:models}, we briefly describe and motivate the cosmological models investigated in this work, while we compare them in Sect.~\\ref{sect:res} with the $\\Lambda$CDM model. We present our conclusions in Sect.~\\ref{sect:conc}. In the appendix \\ref{sect:app}, we discuss why the previous theoretical estimations of $\\delta_\\mathrm{c}$ for the EDE models obtained by \\cite{Bartelmann2006} and \\cite{Sadeh2007} differ from the results obtained in this work and by \\cite{Francis2009a}. ",
        "conclusions": "\\label{sect:conc}  In this work, we have generalised the non-linear equation governing the evolution of matter overdensities to cosmological models containing fluids with an arbitrary equation-of-state parameter. Specifying $w=0$, we recover the well-known equation for structure formation in matter-dominated universes. By means of the non-linear evolution equation, we determine the appropriate initial conditions used to solve the linear equation and compute the linear threshold for collapse, $\\delta_{\\mathrm{c}}$. We point out that the derivation of Eq.~(\\ref{eqn:nldeq}) is very general and can be extended to very broad classes of cosmological models, once the appropriate continuity, Euler, Poisson and background equations are provided.  In this work, we considered exclusively non-clustering dark energy models, in which the only clustering component is the dark matter. Our goal was to study a whole catalogue of dynamical dark-energy models, thereby summarising results partly obtained elsewhere, and to clarify discrepant results on early dark-energy models. We stuck to the common assumption that the sound speed in the dark energy is given by the equation of state. There are more general scenarios allowing perturbations also in the dark-energy component, and this might lead to important differences compared to the case of non-clustering dark energy \\citep{Abramo2007}. If the dark-energy component is not homogeneous, one might wonder what happens to the number counts if the equation of state of the dark-energy fluid changes within the collapsing sphere. Then, the equation of state will depend on the actual overdensity of the dark energy, and the model will acquire an additional degree of freedom, parametrised by the effective sound speed defined by $c_{\\mathrm{s,eff}}^2=\\delta P/\\delta\\rho$. Theoretical predictions for the equation of state inside the collapsing sphere are given by \\cite{Abramo2008}, and a determination of the effective sound speed in Sunyaev-Zel'dovich and weak-lensing surveys is carried out in \\cite{Abramo2009.1}. The authors found that a negative pressure perturbation to the dark energy fluid may have a substantial effect on the number counts as shown in their Fig.~1. A negative effective sound speed $c_s^2$ can lower the value of $\\delta_\\mathrm{c}$ to $\\approx1.5\\ldots1.55$, giving a substantial boost to structure formation. Volume effects will still be comparable with what we found here. It may then be possible to descriminate this class of models compared to those studied here.  Despite the fact that the differential mass functions are very similar because of the small differences in $\\delta_{\\mathrm{c}}$, we still expect a significant difference for the total number of objects in a given volume above a given mass threshold. For low minimum mass, the volume effects dominate and we expect a lower overall number of objects, while the mass function dominates over the volume effect at the high-mass tail, and we still expect more haloes. Unfortunately, as shown in the lower panels of Fig.~\\ref{fig:mf}, the major contributions come from redshifts above $z=1$, where fewer objects are expected and where observations are more difficult.  We found that all models studied here show differences at the per-cent level compared to the standard $\\Lambda$CDM model. We also argue that in the framework of general relativity, it may be possible to have a more pronounced impact on non-linear structure formation if the equation-of-state parameter is discontinuous or if modifications to general relativity are involved. This will be the subject of future work.  We thus conclude that, at least for the wide class of models studied here and due to the current status of observations, it is quite difficult to use number counts to discriminate between different dark-energy models. A discrimination may still be possible to eventually based on geometrical tests where volume effects are relevant."
    },
    "1005/1005.0469_arXiv.txt": {
        "abstract": "We report on the identification of a third, new ultraluminous X-ray source (ULX) in the starburst galaxy M82. Previously, the source was observed at fluxes consistent with the high state of Galactic black hole binaries (BHBs).  We observe fluxes up to $(6.5 \\pm 0.3) \\times 10^{39}$~ergs~s$^{-1}$ in the ultraluminous regime. When the source is not in the the low/hard state, spectral fitting using a multicolor disk model shows that the disk luminosity varies as the disk inner temperature raised to the power $4.8 \\pm 0.9$, consistent with the behavior of Galactic BHBs in the thermal dominant state. Fitting the spectrum with a multicolor disk model with general relativistic corrections suggests that the source harbors a rapidly spinning black hole with a mass less than 100 solar masses. A soft excess was found in the source spectrum that could be blackbody emission from a photosphere created by a massive outflow. The source also showed soft dips during a flare. ",
        "introduction": "Ultraluminous X-ray sources (ULXs) are nonnuclear, point-like X-ray sources in external galaxies with an apparent luminosity over the Eddington limit of a stellar-mass black hole, which is typically $3 \\times 10^{39}$ ergs~s$^{-1}$ for a 20 $M_\\sun$ black hole. Those with variability on time scales from minutes to months are accreting compact objects. The nature of ULXs is still unclear. If the emission is (roughly) isotropic and under the Eddington limit, they may harbor black holes up to $10^3$~$M_\\sun$, belonging to the missing population of intermediate-mass black holes \\citep[IMBHs;][]{col99,mak00,kaa01,far09}. However, they could also be stellar-mass black holes, with beamed or super-Eddington emission \\citep{kin01,kor02,wat01,beg02}.  ULXs with relatively low luminosities ($\\sim 3-6 \\times 10^{39}$~ergs~s$^{-1}$) may be more likely to be stellar-mass black holes of about 10~$M_\\sun$ like Galactic BHBs or up to a couple tens of $M_\\sun$ like IC 10 X-1 \\citep{pre07,sil08}. Galactic BHBs have been observed in a few cases with luminosities up to $6 \\times 10^{39}$ ergs~s$^{-1}$ \\citep{mcc06}, although there are large uncertainties on their distances. \\citet{gla09} proposed that most ULXs are stellar-mass black holes in an `ultraluminous state' with super-Eddington accretion based on analysis of high quality X-ray spectra. A major difference between the two populations is that Galactic BHBs spend most of their lifetime in the quiescent or low state, while most ULXs have persistent emission and have been active since their discovery at time scales up to decades in some cases. The non-transient behavior of the majority of ULXs implies that they do not contain black holes significantly more massive than $10^2$~$M_\\sun$ \\citep{kal04}. Despite these hypotheses, the connection between ULXs and Galactic BHBs has not been well addressed. If they contain black holes of similar mass, we may expect to see transitions from the normal states to the ultraluminous state in an individual source.  Multicolor blackbody emission from an optically thick accretion disk could shed light on the nature of the central black hole. At a fixed fraction of Eddington luminosity, the disk inner temperature scales with the black hole mass to the $-1/4$ power \\citep{mak00}, indicating that accretion disks around IMBHs are cooler and brighter than those around stellar-mass black holes. Soft excesses are detected in ULXs and can be modeled as cool disk emission from IMBHs \\citep{kaa03}. However, the cool disk model is not the unique interpretation of the soft excess \\citep[for a brief review see the introduction of][]{fen09}. Massive outflows due to near or super-Eddington accretion may produce blackbody emission at the similar temperature \\citep{kin03,kin04,beg06,pou07}.  The disk model can be tested by determining if the disk luminosity varies with the 4th power of the inner disk temperature, i.e.\\ $L_{\\rm disk} \\propto T_{\\rm in}^4$.  This relation has been robustly verified for stellar-mass black holes when the disk emission dominates the X-ray spectrum \\citep{gie04}. Although the X-ray spectra of some ULXs can be modeled by dominant disk emission \\citep{sto06,win06}, the observed luminosity variation and number of observations have been inadequate for proper testing of the $L_{\\rm disk} \\propto T_{\\rm in}^4$ relation. A possible $L \\propto T_{\\rm in}^4$ relation was found in NGC 5204 X-1, but the disk emission was not the dominant component in the spectrum and a large correction for Comptonization was required \\citep{fen09}.  M82 is a nearby starburst galaxy in which two ULXs have been found \\citep{mat01,kaa01,fen07b,kon07}. It has been extensively studied in X-rays with the Chandra X-ray observatory . In our recent Chandra observation on 2008 October 4, the X-ray source CXOM82 J095546.6+694041 was 5-10 times brighter than its normal flux level and was identified as the third ULX in M82. According to the naming convention suggested by \\citet{kaa01}, we refer to this source as X37.8+54. In this paper, we investigated its spectral and timing behavior using all available Chandra observations, and discussed its possible nature. The distance to the host galaxy is adopted as 3.63 Mpc \\citep{fre94}. ",
        "conclusions": "We identify X37.8+54 as the third ULX in the starburst galaxy M82 based on its brightening in a recent Chandra observation to an observed flux 5-10 times higher than previously observed. The peak isotropic luminosity in the 0.3-8 keV band corrected for interstellar absorption is $(6.5 \\pm 0.3) \\times 10^{39}$~ergs~s$^{-1}$ inferred from fitting with a power-law model.  The source was often found at relatively low luminosity, near $(1-2) \\times 10^{39}$~ergs~s$^{-1}$. At the lowest luminosities, $\\lesssim 1.5 \\times 10^{39}$~ergs~s$^{-1}$, the source appears to have hard spectra with the photon index less than 2.3.  As the luminosity increases above this value, it shows dramatic spectral variation with the power-law photon index $\\Gamma$ varying from 1.9 to 3.0.  This suggests a change in spectral state, and, indeed, most of the higher luminosity spectra are better fitted with an MCD spectral model.  The lower panel of Figure~\\ref{fig:mcd} shows that when observations in the low/hard state are excluded, the remaining observations exhibit a robust correlation between $L_{\\rm disk}$ and $T_{\\rm in}$. The correlation coefficient between $\\log L_{\\rm disk}$ and $\\log T_{\\rm in}$ is 0.997 with a chance probability of $1.7\\times 10^{-4}$. Fitting the data to a power-law relation, we obtained a power-law slope of $4.8 \\pm 0.9$, consistent with the 4th-power relation predicted for the MCD model in the thermal dominant state. We note that the disk luminosity varies by about an order of magnitude.  Thus, the data provide an adequate test of this relation. We conclude that X37.8+54 exhibits two spectral states: a low/hard state and a thermal dominant state.  Following the recipe in \\citet{mak00} and adopting canonical correction coefficients with $\\kappa=1.7$ and $\\xi=0.412$ ,where $\\kappa$ is the hardening correction factor and $\\xi$ is to correct for the radius of maximum temperature, the disk luminosity can be expressed as a function of the black hole mass ($m$ in solar mass) and the disk inner temperature ($T_{\\rm in}$ in keV) as $L_{\\rm disk} = 7.2\\times 10^{36} \\alpha^2 m^2 T_{\\rm in}^4$~ergs~s$^{-1}$, where $\\alpha$ is related to the black hole spin with $\\alpha=1$ for a non-spinning black hole and $\\alpha=1/6$ for maximal spinning. Using a 4th power-law relation to fit the data points in the bottom panel of Figure~\\ref{fig:mcd}, we obtain $L_{\\rm disk} = (3.7 \\pm 0.3)T_{\\rm in}^{4}/\\cos i \\; 10^{38}$~ergs~s$^{-1}$. Thus, we have $m = (7.2 \\pm 0.3)/(\\alpha\\sqrt{\\cos i})$, suggesting the ULX contains a black hole of about 10 solar masses if it is a Schwartzchild black hole with a face-on disk, or a few tens solar masses if it is fast spinning or viewed at a high inclination.  The best-fit compact object mass is $36 \\pm 20$~$M_\\sun$, derived from fitting with the fully relativistic disk model, indicative of a stellar black hole slightly more massive than Galactic ones. We also tried different hardening correction factors from 1.5 to 1.9. The fitted mass varies by 20\\% with respect to that from $\\kappa=1.7$. Interestingly, the specific angular momentum converges to its maximally allowed value, $a_\\ast = 0.9986$ with a lower limit of 0.80, suggesting the source contains a fast spinning black hole. The mass accretion rate in Eddington unit is about 1.0, which is larger than that found in the thermal dominant state of Galactic BHBs of no more than 0.6 \\citep{mcc06}, but is under the theoretical upper limit of about 10 for a stellar-mass black hole with a radiation dominated accretion disk \\citep{beg02}. \\citet{hui08} investigated the spectra of six disk-dominated ULXs and found five sources had best-fit masses between 23 and 73 $M_\\sun$. The three ULXs that have a derived mass $\\gtrsim$ $25~M_\\sun$ at a confidence $\\gtrsim 99.9\\%$ also show rapid spin, which is quite similar to X37.8+54. The best-fit black hole mass of X37.8+54 is also similar with that of IC~10~X-1 \\citep{pre07,sil08}. It has been suggested that IC~10~X-1 has a persistently high accretion rate \\citep{bar08}; a rapidly spinning black hole is naturally expected since accretion would bring a fair amount of angular momentum to the compact object.  We conclude that stellar Kerr black holes with a mass of several tens $M_\\sun$ could appear as ULXs with a thermal spectrum of a few keV in the 1-10 keV band, as having been suggested by \\citet{mak00}. Such emission could also appear as a hard power-law spectrum if the statistics is inadequate to detect the spectral curvature. These black holes are more massive than Galactic black holes, which are mostly around 10 $M_\\sun$, and may be formed by the collapse of massive stars in a low-metallicity environment and account for a portion of ULXs \\citep{zam09}.  Similar sources, Suzaku J1305-4931 in NGC 4945 and NGC 253 X-2, have been reported in the literature, which show a $L_{\\rm disk} \\propto T_{\\rm in}^4$ correlation as their luminosities varied by a factor of about 2-3 \\citep{iso08,kaj09}. From surveys of bright ULXs in nearby galaxies, many hard ULXs were found and interpreted as IMBHs in the low/hard state \\citep{fen09,kaa09}. For those sources, most of them present a constant spectral index or hardness despite large variation of the luminosity. Therefore, those sources may have a different nature from X37.8+54.  It has been suggested that ULXs are a special case (face-on) of the Galactic microquasar SS~433, where super-Eddington accretion drives a powerful wind off the accretion disk \\citep{fab01,beg06}. The wind could be optically thick and produce a blackbody continuum \\citep{kin03,kin04}. This may be responsible for the soft excess observed in many ULXs with a temperature of about 0.1-0.3 keV \\citep{sto06,pou07,fen07a}. When fitting spectra not from the low/hard state with an MCD component, all observations except 10027 need an additional blackbody component at a significance of at least 3 $\\sigma$. No obvious correlation was found between the luminosity and temperature of the soft blackbody.  However, the uncertainties are large and this result is not very constraining. The calculated surface area of the soft blackbody emission is about $10^{17}$~cm$^2$ for all observations, corresponding to a spheric radius of $10^8$~cm, about $70R_{\\rm g}$ for a 10 solar mass black hole, where $R_{\\rm g}$ is the gravitational radius equal to $GM/c^2$.  The absorption column density varies significantly between observations no matter what models are used. Such large variation is likely internal to the binary, implying the outflow is inhomogeneous and time variable.  We note that the soft blackbody appears at energies where the absorption is important. Thus caution is warranted when interpreting it.  Interesting timing behavior is seen in one of the major outbursts caught by the observations 5644 and 6361 (Figure~\\ref{fig:lc}). The source shows soft dips on time scales of several ks. Soft dips could be caused by absorption of the inner accretion disk, like in 4U~1630$-$47 \\citep{tom98}. This interpretation is also consistent with our spectral analysis results in the penultimate paragraph of Section 2. We note that the soft blackbody component is assumed not to suffer extra absorption. A dense massive outflow from the disk could be the origin of the extra absorption and also the source of the soft blackbody. Other mechanisms could also produce soft dips, e.g., the possible ejection of the inner part of accretion disk into relativistic jets in GRS~1915$+$105 \\citep{nai01}. The time scale of ks of the dip in X37.8+54 is longer than that in 4U~1630$-$47 and in GRS~1915$+$105 by about an order of magnitude. More detailed study of the soft dip requires data with better quality.  To summarize, X37.8+54 provides us a good candidate of a stellar Kerr black hole being ultraluminous with a near Eddington accretion rate. This ULX could be representative of a class of ULXs with relatively low luminosity, i.e.\\ a few $10^{39}$~ergs~s$^{-1}$. The luminosity of X37.8+54 is usually not high enough to be classified as a ULX, while in the outburst it truly enters the ultralumious regime. The X-ray spectrum at peak luminosity is dominated by a disk component, which is unlike many brighter ULXs at their maximum luminosity that they enter the `ultraluminous state' characterized by a strong Comptonizing component \\citep{gla09}. Therefore, this source may be in between the low luminosity tail of ULXs and high luminosity tail of Galactic BHBs. We note that the hard state of X37.8+54 is more luminous than the canonical hard state found in Galactic BHBs. This is consistent with the correlation found in outbursts of BHBs that the peak luminosity of the hard state is scaled with the peak luminosity of the thermal state in the following outburst \\citep{yu09}."
    },
    "1005/1005.3760_arXiv.txt": {
        "abstract": "We present a two-fluid magnetohydrodynamics (MHD) model of quasi-stationary, two-dimensional magnetic reconnection in an incompressible plasma composed of electrons and ions. We find two distinct regimes of slow and fast reconnection. The presence of these two regimes can provide a possible explanation for the initial slow build up and subsequent rapid release of magnetic energy frequently observed in cosmic and laboratory plasmas. ",
        "introduction": "Introduction } Magnetic reconnection is the physical process by means of which magnetic field lines join one another and rearrange their topology. Magnetic reconnection is believed to be the mechanism by which magnetic energy is converted into kinetic and thermal energy in the solar atmosphere, the Earth's magnetosphere, and in laboratory plasmas~\\cite{biskamp_2000,priest_2000,bc_low_2003,kulsrud_2005,drake_2006, yamada_2009,zweibel_2009}. Many reconnection related physical phenomena observed in cosmic and laboratory plasmas exhibit a two-stage behavior. During the first stage, magnetic energy is slowly built up and stored in the system with relatively little reconnection occurring. The second stage is characterized by a sudden and rapid release of the accumulated magnetic energy due to a fast reconnection process. For example, a solar flare is powered by a sudden (on timescale ranging from minutes to tens of minutes) release of magnetic energy stored in the upper solar atmosphere~\\cite{kulsrud_2005}. Because the value of the Spitzer electrical resistivity is very low in hot plasmas, magnetic energy release rates predicted by a simple single-fluid MHD description of magnetic reconnection are much slower than the rates observed during fast reconnection events in astrophysical and laboratory plasmas~\\cite{biskamp_2000,bc_low_2003,kulsrud_2005,drake_2006, yamada_2009,zweibel_2009}. One of the most promising solutions of this discrepancy is the two-fluid MHD theoretical approach to magnetic reconnection~\\cite[and references therein]{biskamp_2000, kulsrud_2005,drake_2006,yamada_2009,zweibel_2009}. Recently a model of two fluid reconnection in a electron-proton plasma was presented in~\\cite{malyshkin_2009}. In this paper, we consider a more general case of two-fluid reconnection in electron-ion and electron-positron plasmas, and we present derivations in detail. In the discussion section, we also argue that the slow and fast reconnection regimes predicted by our model, can provide a possible explanation for the observed two-stage reconnection behavior. ",
        "conclusions": "Discussion }  The solution for two-fluid reconnection is summarized in table~\\ref{TABLE_SOLUTION}. This table includes solution formulas for three reconnection regimes: the slow Sweet-Parker reconnection regime, the transitional Hall reconnection regime, and the fast collisionless reconnection regime. The reconnection rates for these three regimes are respectively shown by the solid, dotted and dashed lines in figure~\\ref{SOLUTION_PICTURE}.  It is well known that resistivity $\\eta$ can be considerably enhanced by current-driven plasma instabilities~\\cite{yamada_2009,zweibel_2009,kulsrud_2001}. Because the collisionless reconnection rate $E_z\\approx\\eta B_{ext}/d_e$ is proportional to the resistivity [see equation~(\\ref{Ez_FAST})], this rate can increase significantly as well. As a result, we propose the following possible theoretical explanation for the two-stage reconnection behavior (fast and slow) that is frequently observed in cosmic and laboratory plasma systems undergoing reconnection processes.   \\begin{figure}[t] \\vspace{5.0truecm} \\special{psfile=solution_picture.eps voffset=5 hoffset=100 vscale=40 hscale=40 angle=0} \\caption{Schematic plot of the reconnection rate $E_z$ versus the Lundquist number $S$ in the slow Sweet-Parker (solid line), transitional Hall (dotted line), and fast collisionless (dashed line) reconnection regimes. } \\label{SOLUTION_PICTURE} \\end{figure}   During the first stage, such a system is in the very slow Sweet-Parker reconnection regime, during which magnetic energy is slowly built up and stored in the system. The magnetic energy and electric currents build up, the field strength increases and the resistivity decreases~\\cite{uzdensky_2007}. As a result, the Lundquist number $S$ increases and the system moves to the right along the solid line in figure~\\ref{SOLUTION_PICTURE}.  When the Lundquist number $S$ becomes comparable to $L_{ext}^2/d_i^2$ and the thickness of the current layer $L_{ext}/S^{1/2}$ becomes comparable to $d_i$, the system reaches point~A in figure~\\ref{SOLUTION_PICTURE}. Next the system goes into the transitional Hall reconnection regime and quickly moves up along the vertical dotted line in figure~\\ref{SOLUTION_PICTURE}. During this transition, the length of the electron layer shrinks from $\\approx L_{ext}$ to $\\approx(d_e/d_i)L_{ext}$, the electron layer thickness decreases from $\\approx d_i$ to $\\approx d_e$, and both the electric current and the reconnection rate increase by a factor $\\approx d_i/d_e\\gg 1$. The system ends up in the fast collisionless reconnection regime at point~B in figure~\\ref{SOLUTION_PICTURE}.  Because of the considerable increase in the electric current during the Hall reconnection transition from point~A to point~B, plasma instabilities develop, and, consequently, resistivity $\\eta$ becomes anomalous and rises in value. As a result, the reconnection rate $E_z\\approx\\eta B_{ext}/d_e$ increases, the Lundquist number $S=V_AL_{ext}/\\eta$ and electron layer length $L\\approx V_Ad_ed_i/\\eta$ decrease, and the system moves from point~B to the left along the dashed line in figure~\\ref{SOLUTION_PICTURE}. The system enters the second stage characterized by a rapid release of the accumulated magnetic energy. Even though our theoretical model is stationary, assumes constant resistivity and cannot describe this stage in detail, the physical mechanism of slow and fast reconnection outlined above is self-consistent and may take place in nature."
    },
    "1005/1005.1286_arXiv.txt": {
        "abstract": "We explore model-independent collider constraints on light Majorana dark matter particles. We find that colliders provide a complementary probe of WIMPs to direct detection, and give the strongest current constraints on light DM particles. Collider experiments  can access interactions not probed by direct detection searches, and outperform direct detection experiments by about an order of magnitude for certain operators in a large part of parameter space. For operators which are suppresssed at low momentum transfer, collider searches have already placed constraints on such operators limiting their use as an explanation for DAMA. ",
        "introduction": "Recently, there has been much interest in light (order $\\sim$~GeV) mass dark matter \\cite{Fitzpatrick:2010em,Kuflik:2010ah,Feldman:2010ke,Chang:2010yk,Essig:2010ye}. This interest is partly spurred by the fact that the DAMA signal of annual modulation \\cite{Bernabei:2010mq} may be understood as consistent with null results reported by CDMS \\cite{Ahmed:2009zw} and Xenon 10 \\cite{Angle:2007uj} if the dark matter is a weakly interacting massive particle (WIMP) of mass  $\\lesssim 10$~GeV \\cite{Petriello:2008jj}. Further excitement is motivated by the signal reported by CoGeNT, which favors a WIMP in the same mass range \\cite{Aalseth:2010vx} as DAMA with moderate channeling (however, unpublished data from 5 towers of CDMS Si detectors \\cite{CDMS-Si} provides some tension, see \\cite{Chang:2010yk}).  A WIMP which is relevant for direct detection experiments necessarily has substantial coupling to nucleons, and thus can be produced in high energy particle physics experiments such as the Tevatron and Large Hadron Collider (LHC). In particular, light WIMP states can be produced with very large rates. These WIMPs escape undetected, and hence the most promising signals involve missing energy from a pair of WIMPs recoiling against Standard Model (SM) radiation from the initial state quarks/gluons \\cite{Birkedal:2004xn,Cao:2009uw,Beltran:2010ww}.  While such searches are complicated % by large SM backgrounds producing missing energy, % we will find that colliders can provide stringent restrictions on the parameter space of light dark matter models. Colliders can also access interactions which are irrelevant for direct detection (either because they lead to vanishing matrix elements in non-relativistic nucleon states or are suppressed at low momentum transfer).  In this article, we explore the bounds colliders can place on a light Majorana fermion WIMP, which we assume interacts with the SM largely through higher dimensional operators. By exploring the complete set of leading operators, we arrive at a model-independent picture (up to our assumptions) of WIMP interactions with SM particles in the case where the WIMP is somewhat lighter than any other particles in the dark sector. We show that colliders can outperform direct detection searches significantly over a large area of parameter space. ",
        "conclusions": "We have derived new constraints on generic Majorana DM models based on null search results for monojets  at the Tevatron, and explored corresponding discovery reaches at the LHC. Our bounds cover regions of parameter space which were previously not constrained by experimental efforts and strongly constrain some kinds of low mass WIMPs as an explanation for the DAMA and CoGeNT signals. In particular, we have derived constraints on the direct detection scattering of light Majorana WIMPs which are significantly stronger than experimental bounds (and near-future prospects) for spin-dependent scattering.  Colliders are particularly good experiments for testing DM models which are suppressed at small momentum transfer, whether the suppression is kinematic in nature as in models of light DM, or if the momentum dependence is inherent in the induced operator itself, as is the case with momentum-dependent DM\\,\\cite{Chang:2009}. It would be interesting to study the collider constraints on these models in detail.\\\\   \\noindent {\\em Note added:}  During the final stages of preparing this manuscript, the Xenon 100 collaboration released data constraining the low WIMP mass region of parameter space \\cite{Aprile:2010um}."
    },
    "1005/1005.1848_arXiv.txt": {
        "abstract": "{ With appropriate spatial resolution, images of spiral galaxies in thermal infrared ($\\sim$10\\,\\mic\\ and beyond) often reveal a bright central component, distinct from the stellar bulge, superimposed on a disk with prominent spiral arms. {\\em ISO} and {\\em Spitzer} studies have shown that much of the scatter in the mid-infrared colors of spiral {\\revision galaxies is related to changes} in the relative importance of these two components, rather than to other modifications, such as  the morphological type or star formation rate, that affect the properties of the galaxy as a whole. With the {\\em Herschel} imaging capability from 70 to 500\\,\\mic, we revisit this two-component approach at longer wavelengths, to see if it still provides a working description of the brightness distribution of galaxies, {\\revision and to determine its implications on the interpretation of global far-infrared properties of galaxies.}  We quantify the luminosity of the central component by both a decomposition of the radial surface brightness profile and a direct extraction in 2D. We find the central component contribution is variable {\\revision within the three galaxies in our sample}, possibly connected more directly to the presence of a bar than to the morphological type. The central component's relative contribution is at its maximum in the mid-infrared range and drops around 160\\,\\mic\\ to reach a constant value beyond 200\\,\\mic. The central component contains a greater fraction of hot dust than the disk component, and while the colors of the central components {\\revision are} scattered, colors of the disk components are more homogenous from one galaxy to the next. } ",
        "introduction": "The central region of spiral galaxies hosts a number of processes that have little or no counterpart in the disks. A nuclear starburst, triggered by bar instabilities and associated or not with an active galactic nucleus (AGN) can often be found. The stellar density is also enhanced with respect to the disk, leading to an increase in the interstellar radiation field (ISRF) intensity. It is thus no surprise that in the mid-infrared (MIR), spiral galaxies can often be seen as the sum of a central component, hereafter referred to as the core component as distinct from the bulge, and a disk component \\citep[see e.g.][]{bendo07}. Extraction of the core component from MIR brightness profiles have revealed that it can contribute up to 90\\% of the total MIR luminosity and that the distribution of core/total luminosity ratio is evenly spread between 0 and 1 \\citep{rousselbar}. This study also showed that much of the scatter in the MIR {\\em global} properties of spiral galaxies came from changes in the relative contribution of the core component and of its spectral energy distribution (SED), while disk components formed a much more homogenous group. In short, the scatter affecting the {\\em global} MIR colors of galaxies is caused by a very {\\em local} effect, i.e., the importance of the core component, rather than by a change in any single property of galaxies such as the ISRF throughout the disk. Taking this into account, \\citet{rousselsfr} showed that the MIR luminosity of galaxies could be turned into an accurate tracer of the star-formation rate, provided it was corrected for the core contribution. More recently, \\citet{chanial} extended this analysis by showing that the dispersion in the correlation between the infrared luminosity of galaxies and the dust temperature could be reduced by considering a compactness parameter, again related to the relative contribution of the core and disk components.  Care must therefore be taken when extrapolating locally established relations, such as the correlation of star-formation tracers in galactic disks, to global studies where galaxies are no longer resolved. One first needs to verify that the total luminosity of galaxies is indeed dominated by the component where these correlations have been established, or to design methods to correct for the existence of multiple contributors to the total luminosity. With its unprecedented sensitivity and spatial resolution, {\\em Herschel} \\citep{pilbratt} now allows us to carry this very investigation over the full thermal infrared peak, by combining the PACS instrument \\citep{poglitsch}, delivering spatial resolutions of $5\\farcs5$, $6\\farcs7$, and $11''$ (FWHM) at 70, 100, and 160\\,\\mic, and the SPIRE instrument \\citep{griffin} delivering spatial resolutions of $18\\farcs1$, $25\\farcs2$, and $36\\farcs9$ (FWHM) in its three photometric bands at 250, 350, and 500\\,\\mic. ",
        "conclusions": "\\label{sec:conc}  In this small set of galaxies, the core fractions are small ($\\leq{20}\\%$), therefore the correction to the full SED of the galaxies introduced by accounting for the different core properties would be minimum. Nevertheless we already observe both a significant variation in the core component contribution from one galaxy to the next and a systematic decrease with wavelength. We also find that the galactic disks studied here show remarkable color homogeneity over the 24-500\\,\\mic\\ range, while the core components are notably hotter as shown by their higher [70]/[160] color value. This opens the possibilities that (i) across the {\\em Herschel} wavelength range, the global SED of a galaxy may be influenced as much by the relative importance of the core and disk components as by a global property, such as the star formation rate or the morphological type, and (ii) that this may depend on the observed wavelength. We shall thus examine this in larger samples such as the HRS set of resolved spiral galaxies."
    },
    "1005/1005.4189_arXiv.txt": {
        "abstract": "We study the emission of X-rays from lobes of FR-II radio galaxies by inverse Compton scattering of microwave background photons. Using a simple model that takes into account injection of relativistic electrons, their energy losses through adiabatic expansion, synchrotron and inverse Compton emission, and also the stopping of the jet after a certain time, we study the evolution of the total X-ray power, the surface brightness, angular size of the X-ray bright region and the X-ray photon index, as functions of time and cocoon size, and compare the predictions with observations. We find that the radio power drops rapidly after the stopping of the jet, with a shorter time-scale than the X-ray power. The X-ray spectrum initially hardens until the jet stops because the steepening of electron spectrum is mitigated by the injection of fresh particles, for electrons with $\\gamma \\ge 10^3$. This happens because of the concurrence of two times scales, that of the typical jet lifetimes and cooling due to inverse Compton scattering ($\\sim 10^{7\\hbox{--}8}$ yr), of electrons responsible for scattering CMB photons into keV range photons (with $\\gamma \\sim \\sqrt{1 \\, {\\rm keV}/ kT_{CMB}}$). Another finding is that the ratio of the X-ray to radio power is a robust parameter that varies mostly with redshift and ambient density, but is weakly dependent on other parameters. We also determine the time-averaged ratio of X-ray to radio luminosities (at 1 keV and 151 MHz) and find that it scales with redshift as $\\propto (1+z)^{3.8}$, for typical values of parameters. We then estimate the X-ray luminosity function of FR-II radio galaxies and estimate the number of these diffuse X-ray bright objects above a flux limit of $\\sim 3 \\times 10^{-16}$ erg cm$^{-2}$ s$^{-1}$ to be $\\sim 25$ deg$^{-2}$. ",
        "introduction": "A number of radio galaxies have been observed to emit diffuse X-rays in the region between the nucleus and radio hot spots in recent years. This emission has been interpreted as inverse-Compton scattering of the cosmic microwave background (CMB) photons by relativistic electrons in the radio lobes that have lost most of their energy and their Lorentz factor have come down to $\\gamma \\sim 10^3$ (Fabian \\etal 2003; Blundell \\etal 2006; Erlund \\etal 2008; Johnson \\etal 2007; Croston \\etal 2005; Fabian \\etal 2009; Isobe \\etal 2009). That relativistic electrons in radio galaxy cocoons could upscatter CMB photons to X-rays has long been anticipated since the discovery of CMB (e.g., Felten \\& Rees 1967).   Since the increase in the CMB energy density with redshift ($\\propto (1+z)^4$) compensates for the cosmological surface brightness dimming, this extended emission has been seen at both low and high redshift, and has been used as a probe of the relativistic plasma in the lobe.  Combined with radio observations, this emission provides constraints on the magnetic field and on the energy distribution among the relativistic particles, or both. While electrons with $\\gamma \\ge10^4$ to emit GHz synchrotron radiation (for typical magnetic fields of a few $\\mu$G), it requires $\\gamma \\sim 10^3$ to upscatter the CMB photons to the observed $1\\hbox{--}10$ keV range in X-rays.  The simultaneous radiation in two different wavelengths caused by the same parent electron population provides an excellent opportunity to probe the physical nature of the radiating object. Comparing the X-ray and radio emission from lobes, Blundell \\etal (2006) showed that the injected electron energy distribution have a low-energy turn-over at $\\gamma _{min} \\ge 10^3$, since the X-ray bright regions did not coincide with the radio hot spots. The comparison between X-ray and radio studies also shed light on the magnetic field. Croston \\etal (2005) argued that the inferred magnetic field strength is close to the equipartition value (lying between $0.3\\hbox{--}1.3$ times the equipartition value).  In this paper, we use a variant of a model of emission from FRII radio galaxies advocated by Kaiser \\& Alexander (1997) (hereafter KA97) and Kaiser, Dennett-Thorpe \\& Alexander (1997) (hereafter KDA97) to study the evolution of X-ray power and surface brightness with time, and as functions of the radio lobe size, ambient density, redshift and the jet power. We do not assume self-similarity of the evolution of the radio lobe assumed by KA97 and KDA97, in order to study the effect of the cessation of jet activity on the X-ray surface brightness.   The evolution of radio lobes after the stopping of jets has been discussed by Komissarov \\& Gubanov (1994), although they did not address the soft X-ray emission by old electrons. Kaiser \\& Cotter (2002) have also discussed the evolution of radio galaxies in connection with the observations of radio relics, and Reynolds \\etal ( 2002) have numerically studied the expansion of 'dead' radio galaxies.   We first discuss the KA97 and KDA97 model and the changes introduced in the model in the present paper, and then discuss the results of our calculations. ",
        "conclusions": "We have used an evolving model of FR-II radio galaxy cocoons that takes into account the stopping of the jet after a jet lifetime of $t_j$, and we have studied the X-ray property of cocoons from inverse Compton scattering of CMB photons. We have shown that the the X-ray power (in the 1-5 keV band) and the X-ray surface brightness decrease rapidly after the stopping of the jet, as well as the radio power. At the same time, the X-ray photon index is found first decrease to some extent while the jet is active. This occurs because at the important energy scale for inverse Compton scattering of CMB photons to keV range ($\\gamma \\sim 10^3$), two relevant timescales become comparable: the jet lifetime and that of cooling due to inverse Compton ($10^{7\\hbox{--}8}$ yr). After the stopping of the jet, the photon index increases, making the spectrum softer. We find that typical radio galaxy parameters produce an X-ray flux (1-5 keV band) of order $1\\hbox{--}10$ $\\mu$-Crab, being emitted from regions with sizes of order of tens of arcsec.  We have also shown that the ratio of X-ray to radio power changes little with jet power and lower cut-off in the electron energy distribution, and it changes mostly with redshift and ambient density. We have then determined the time-averaged ratio of X-ray to radio power and studied its scaling with redshift, which we found to scale as $\\propto (1+z)^{3.8}$, almost independent of the injected electron energy spectrum, contrary to the expectations from a population of electrons with a single power law distribution of energy.  Using these scaling relations, we estimated the X-ray luminosity function of FR-II radio galaxies at various redshifts, and estimated the number of such diffuse X-ray sources to be $\\sim 25$ deg$^{-2}$ above an X-ray flux limit of $\\sim 3 \\times 10^{-16}$ erg cm$^{-2}$ s$^{-1}$ (compared to $\\sim 167$ deg$^{-2}$ diffuse soft X-ray sources observed in the {\\it Chandra} Deep Field North for a similar flux limit).  \\bigskip I am grateful to Paddy Leahy, the referee, for valuable comments. I also thank Mitchell Begelman for his comments on the paper and Biswajit Paul for discussions."
    },
    "1005/1005.0413_arXiv.txt": {
        "abstract": "We study the dynamical stability of  planetary systems consisting of one  hypothetical terrestrial mass planet ($1~  $ or $10 \\mearth$) and one massive planet ($10 \\mearth - 10 \\mjup$). We consider masses and orbits that cover the range of observed planetary system architectures (including non-zero initial eccentricities), determine the stability limit through N-body simulations, and compare it to the analytic Hill stability boundary. We show that for given masses and orbits of a two planet system, a single parameter, which can be calculated analytically,  describes the Lagrange stability boundary (no ejections or exchanges) but which diverges significantly from the Hill stability boundary. However, we do find that the actual boundary is fractal, and therefore we also identify a second parameter which demarcates the transition from stable to unstable evolution. We show the portions of the habitable zones of $\\rho$ CrB, HD 164922, GJ 674, and HD 7924 which can support a terrestrial planet. These analyses clarify the stability boundaries in exoplanetary systems and demonstrate that, for most exoplanetary systems, numerical simulations of the stability of potentially habitable planets are only necessary over a narrow region of parameter space. Finally we also identify and provide a catalog of known systems which can host terrestrial planets in their habitable zones. ",
        "introduction": "\\label{sec1}   The dynamical stability of extra-solar planetary systems can constrain planet formation models, reveal commonalities among planetary systems and may even be used to infer the existence of unseen companions. Many authors have studied the dynamical  stability of our solar system and extra-solar planetary systems \\citep[see][for example]{Wisdom1982,Laskar1989, RasioFord1996, Chambers1996, LaughlinChambers2001, Gozdziewski2001, Ji2002,BQ04, Ford2005, Jones2006,raymond09}. These investigations have revealed that planetary systems are close to dynamical instability, illuminated the boundaries between stable and unstable configurations, and identified the parameter space that can support additional planets.  From an astrobiological point of view, dynamically stable habitable zones  (HZs)  for terrestrial mass planets ($0.3 \\mearth < M_{p} < 10 \\mearth$) are the most interesting. Classically, the HZ is defined as the circumstellar  region in which a terrestrial mass  planet with favorable atmospheric conditions can sustain liquid water on its surface \\citep[][but see also Barnes et al.(2009)]{Huang1959, Hart1978, Kasting1993, Selsis2007}.   Previous work  \\citep{Jones2001, MT2003, Jones2006, Sandor2007} investigated the orbital stability of Earth-mass planets in the HZ  of systems with a Jupiter-mass companion. In their pioneering work, \\cite{Jones2001} estimated the stability of four known planetary systems in the HZ of their host stars. \\cite{MT2003} considered the dynamical stability of 100 terrestrial mass-planets (modelled as test particles) in the HZs of the then-known 85 extra-solar planetary systems. From their simulations, they generated a tabular list of stable HZs for all observed systems. However, that study did not systematically consider eccentricity, is not generalizable to arbitrary planet masses, and relies on numerical experiments to determine stability. A similar study by \\cite{Jones2006} also examined the stability of  Earth-mass planets in the HZ. Their results indicated that $41 \\%$ of the systems in their sample had ``sustained habitability''. Their simulations were also not generalizable and based on a large set of numerical experiments which assumed the potentially habitable planet was on a circular orbit. Most recently, \\cite{Sandor2007} considered systems consisting of a giant planet with a maximum eccentricity of $0.5$ and a terrestrial planet (modelled as a test particle initially in circular orbit) They used relative Lyapunov indicators  and fast Lyapunov indicators  to identify stable zones and generated a stability catalog, which can be applied to systems with mass-ratios in the range $10^{-4} - 10^{-2}$ between the giant planet and the star. Although this catalog is generalizable to massive planets between a Saturn-mass and $10 \\mjup$, it still assumes the terrestrial planet is on a circular orbit.  These studies made great strides toward a universal definition of HZ stability. However, several aspects of each study could be improved, such as a systematic assessment of the stability of terrestrial planets on eccentric orbits, a method that eliminates the need for computationally expensive numerical experiments, and wide coverage of planetary masses. In this investigation we address each of these points and develop a simple analytic approach that applies to arbitrary configurations of a giant-plus-terrestrial planetary system.  As of March 2010, 376 extra-solar planetary systems have been detected, and the majority (331, $\\approx 88 \\%$) are single planet systems. This opens up the possibility that there may be additional planets not yet detected, in the stable regions of these systems. According to \\cite{wright2007} more than $30 \\%$ of known single planet systems show evidence for additional companions. Furthermore, \\cite{marcy2005} showed that the distribution of observed planets rises steeply towards smaller masses. The analyses of \\cite{wright2007} \\& \\cite{marcy2005} suggests that many systems may have low mass planets\\footnote{ \\cite{Wittenmyer2009} did a comprehensive study of 22 planetary systems using the Hobby-Eberly Telescope (HET; \\cite{Ramsey98}) and found no additional planets, but their study had a  radial velocity (RV) precision of just $10 \\sim 20$ ms$^{-1}$, which can only detect low-mass planets in tight orbits.}. Therefore, maps of stable regions in known planetary systems can aid observers in their quest to discover more planets in known systems.    We consider two definitions of dynamical stability: (1) {\\it Hill stability}: A system is Hill-stable if the ordering of planets is conserved, even if the outer-most planet  escapes to infinity. (2) {\\it Lagrange stability}: In this kind of stability, every planet's motion is bounded, i.e, no planet escapes from the system and exchanges are forbidden. Hill stability for a two-planet, non-resonant system can be described by an analytical expression \\citep{MB1982, Gladman1993}, whereas no analytical criteria are available for Lagrange stability so we investigate it through numerical simulations. Previous studies by \\cite{BG06, BG07}  showed that Hill stability is a reasonable approximation to Lagrange stability in the case of two approximately Jupiter-mass planets. Part of the goal of our present work is to broaden the parameter space considered by \\cite{BG06, BG07}.  In this investigation,  we explore the stability of hypothetical $1 \\mearth$ and $10  \\mearth$ planets in the HZ and in the presence of giant and super-Earth planets.  We consider nonzero initial eccentricities of terrestrial planets and find that a modified version of the Hill stability criterion adequately describes the Lagrange stability boundary. Furthermore, we provide an analytical expression that identifies the Lagrange stability boundary of two-planet, non-resonant systems.  Utilizing these  boundaries, we  provide a catalog of  fractions of HZs that are Lagrange stable for terrestrial mass planets in all the currently known single planet systems. This catalog can help guide observers toward systems that can host terrestrial-size planets in their HZ.  The plan of our paper is as follows: In Section \\ref{sec2}, we  discuss the Hill and Lagrange stability criteria, describe our numerical methods, and present our  model of the HZ. In Section \\ref{sec3}, we present our results and explain how to identify  the Lagrange stability boundary for any system with one $\\ge 10 \\mearth$ planet and one $\\le 10 \\mearth$ planet. In Section \\ref{sec4}, we apply our results to some of the known single planet systems. Finally,  in Section \\ref{sec5}, we summarize the investigation, discuss its importance  for observational programs, and suggest directions for future research. ",
        "conclusions": "\\label{sec5} We have empirically determined the Lagrange stability boundary for a planetary system consisting of one terrestrial mass planet and one massive planet, with initial eccentricities less than 0.6. Our analysis  shows that for two-planet systems with one terrestrial like planet and one more massive planet, Eq.(\\ref{tauearth}) defines Lagrange stable configurations and can be used to identify systems with HZs stable for terrestrial mass planets. Furthermore, in Table \\ref{fractable} we provide a catalog of exoplanets, identifying the fraction of HZ that is Lagrange stable for terrestrial mass planets. A full version of the table is available in the electronic edition of the journal\\footnote{Updates to this catalog is available at \\url{http://gravity.psu.edu/~ravi/planets/}.}.  In order to identify stable configurations for a terrestrial planet, one can calculate a stability boundary (denoted as $\\tau_\\mathrm{s}$ from Eq.(\\ref{tauearth})) for a given system (depending on the  eccentricity and mass of the known planet), and calculate the range of $a$ \\& $e$ that can support a terrestrial planet, as shown in Section 4. For the transitional region between  unstable to  stable ($\\tau_\\mathrm{u} < \\bbc < \\tau_\\mathrm{s}$), a numerical integration should be made. Our results are in general agreement with previous studies \\citep{MT2003}, \\citep{Jones2006} \\& \\citep{Sandor2007}, but crucially our approach does not (usually) require a large suite of N-body integrations to determine stability.  We have only considered two-planet systems, but  the possibility that the star hosts more, currently undetected planets is real and may change the stability boundaries outlined here. However, the presence of additional companions will likely decrease the size of the stable regions shown in this study. Therefore, those systems that have fully unstable HZs from our analysis will likely continue to have unstable HZs as more companions are detected (assuming the mass and orbital parameters of the known planet do not change with these additional discoveries). The discovery of an additional planet outside the HZ that destabilizes the HZ is also an important information.   As more extra-solar planets are discovered, the resources required to follow-up grows.  Furthermore, as  surveys push to lower planet masses, time on large telescopes is required, which is in limited supply. The study of exoplanets seems poised to transition to an era in which systems with the potential to host terrestrial mass planets in HZs will be the focus of surveys. With limited resources, it will be important to identify systems that can actually support a planet in the HZ. The parameter $\\tau_{s}$ can therefore guide observers as they hunt for the grand prize in exoplanet research,  an inhabited planet.  Although the current work focuses on  terrestrial mass planets, the same analysis can be applied to arbitrary configurations that cover all possible orbital parameters. Such a study could represent a significant improvement over the work of \\cite{BG07}. The results presented here show that $\\bbc=1$ is not always the Lagrange stability boundary, as they suggested. An expansion of this research to a wider range of planetary and stellar masses and larger eccentricities could provide an important tool for determining the stability and packing of exoplanetary systems. Moreover, it could reveal an empirical relationship that describes the Lagrange stability boundary for two planet systems. As new planets are discovered in the future, the stability maps presented here will guide future research on the stability of extra-solar planetary systems."
    },
    "1005/1005.2416_arXiv.txt": {
        "abstract": "At the time of recombination, baryons and photons decoupled and the sound speed in the baryonic fluid dropped from relativistic, $\\sim c/\\sqrt3$, to the thermal velocities of the hydrogen atoms, $\\sim 2\\times 10^{-5}c$.  This is less than the relative velocities of baryons and dark matter computed via linear perturbation theory, so we infer that there are supersonic coherent flows of the baryons relative to the underlying potential wells created by the dark matter.  As a result, the advection of small-scale perturbations (near the baryonic Jeans scale) by large-scale velocity flows is important for the formation of the first structures.  This effect involves a quadratic term in the cosmological perturbation theory equations and hence has not been included in studies based on linear perturbation theory.  We show that the relative motion suppresses the abundance of the first bound objects, even if one only investigates dark matter haloes, and leads to qualitative changes in their spatial distribution, such as introducing scale-dependent bias and stochasticity.  We further discuss the possible observable implications of this effect for high-redshift galaxy clustering and reionization. ",
        "introduction": "The early Universe was extremely homogeneous and isotropic, with small adiabatic density perturbations likely seeded during an epoch of inflation \\cite{1981PhRvD..23..347G, 1982PhLB..108..389L, 1982PhRvL..48.1220A}.  The subsequent evolution of the Universe is well-described by a model containing baryons, cold dark matter (CDM), and a cosmological constant ($\\Lambda$).  This inflationary $\\Lambda$CDM paradigm, with only six parameters, is simultaneously consistent with a wide range of cosmological observables \\cite{2007ApJS..170..377S, 2010arXiv1001.4538K}.  One of the key features of the $\\Lambda$CDM scenario is the hierarchical formation of structure: since the matter power spectrum $\\Delta^2(k)$ is an increasing function of wavenumber $k$, the smallest perturbations collapse first, followed by their assembly into larger and larger structures.  The formation of the first structures has recently become a major research area: aside from the intrinsic interest in understanding the first galaxies, these objects are believed to be responsible at least partially for the reionization of the intergalactic medium \\cite{1994MNRAS.269..563F, 1996ApJ...472L..63O}, and they are sensitive to the small-scale power spectrum of the dark matter, which is a powerful probe of dark matter microphysics.  The evolution of density perturbations in the early Universe is generally described using linear perturbation theory, which treats overdensities and velocity fields as small quantities and hence neglects second order terms.  Several previous works have extended this theory down to the post-recombination baryon Jeans scale \\cite{1997PhRvD..56.7566Y, 1998ApJ...501..442Y}. Interest in direct observations of the high-redshift Universe via absorption in the redshifted 21 cm line \\cite{2004PhRvL..92u1301L} has motivated more detailed investigation of the clustering of baryons during the epoch between recombination and reionization \\cite{2005MNRAS.363L..36B, 2007MNRAS.377..667N}, including the entropy and ionization fluctuations in the baryons \\cite{2005MNRAS.362.1047N, 2007PhRvD..76f3001L, 2007PhRvD..76h3005L}.  A deficiency of linear perturbation theory is that it does not describe the collapse of perturbations to form bound haloes, although analytical models such as the Press-Schechter formalism \\cite{1974ApJ...187..425P, 1991ApJ...379..440B} are often used to estimate the halo mass function and clustering. In order to go beyond linear perturbation theory, one may use spherically symmetric Lagrangian hydrodynamic models \\cite{1996ApJ...464..523H} or high-resolution 3D hydrodynamic simulations to follow the infall of baryons into the first haloes \\cite{2002ApJ...564...23B, 2009MNRAS.399..369N}. Since it is inherently nonperturbative this approach can, with incorporation of appropriate chemistry and cooling processes, even be followed all the way to the formation of the first stars \\cite{2002Sci...295...93A, 2005ApJ...628L...5O, 2008Sci...321..669Y}.  The principal purpose of this paper is to point out a new nonlinear effect in the growth of small-scale density perturbations that is active even at $z\\sim 1000$.  The idea is that prior to recombination, the baryons are tightly coupled to the photons resulting in a standing acoustic wave pattern \\cite{1970Ap&SS...7....3S}.  Modern linear perturbation theory treatments including the CDM \\cite{1995ApJ...444..489H, 1995ApJ...455....7M} show a consequent relative velocity of the baryons and CDM since the latter does not suffer Thomson scattering and merely follows geodesics of the cosmic spacetime.  At the time of recombination, the root-mean-square (RMS) relative velocity is 30 km s$^{-1}$, and this is coherent over a scale of several Mpc comoving (the Silk damping scale \\cite{1968ApJ...151..459S}).  When the baryons recombine and are no longer tied to the photons, their sound speed drops to $\\approx6$ km s$^{-1}$, and hence there is a highly supersonic relative velocity between baryons and CDM.  This means that near the baryonic Jeans scale, perturbations in the baryons and CDM are advected relative to each other in less than a Hubble time (and hence less than their growth time).  This effect is investigated herein, and we find that it both suppresses the growth of small-scale structure, and leads to qualitatively new effects in the clustering of the first bound baryonic objects.  The suppression effect does not appear to have been present in previous analyses.  Since it results from the coupling of large-scale and small-scale modes, it is nonlinear and hence not present in linear perturbation theory.  Since the large-scale modes involved are associated with the acoustic oscillations of the photon-baryon fluid, they are not properly modeled by hydrodynamic simulations whose box size is smaller than the acoustic horizon ($\\sim 140$ Mpc comoving).  Understanding of the physical processes that determine the collapse of the first dark matter halos and subsequent accumulation of baryonic matter in those halos is of paramount importance for interpretation of future data on reionization, high-redshift galaxies, and possibly dark matter substructure. In the present paper we introduce the formalism and focus on the key features of the relative velocity effect, leaving a detailed study of various applications for future work.  This paper is organized as follows. In Section~\\ref{sec:PT} we introduce the effect and calculate the Mach number of the relative motion of dark matter and baryonic fluids at the time of recombination. We show that accounting for the relative motion leads to a suppression of the matter power spectrum near the baryon Jeans scale. In Section~\\ref{sec:sim} we compute the abundance and clustering properties of the first haloes taking account of the relative motion.  The treatment is simple (it uses the linear Gaussian random field model for the large-scale density and velocity perturbations in cells of size a few Mpc, and then uses an analytic model to compute the density of small haloes in each cell), but we believe it should capture the qualitative result of the relative motion effect.  We briefly summarize our results and outline possible future work in Section~\\ref{sec:conc}.  The numerical results and plots shown in this paper assume a cosmology with present-day baryon density $\\Omega_{\\rm b,0}=0.044$, CDM density $\\Omega_{\\rm c,0}=0.226$, dark energy density $\\Omega_{\\rm \\Lambda,0}=0.73$, Hubble constant $H_0=71$ km$\\,$s$\\,$Mpc$^{-1}$, and adiabatic primordial perturbations of variance $\\Delta^2_\\zeta=2.42\\times 10^{-9}$. ",
        "conclusions": "\\label{sec:conc}  We have shown that the relative velocity of baryonic and dark matter fluids plays an important role in the formation and evolution of small-scale structure of the early universe.  In light of the increasing interest in reionization, high-redshift galaxy clustering, and dark matter substructure, it is imperative to understand the evolution of small scales structure and all physical effects that contribute to this evolution. Here we discuss the possible implications and next steps in exploring the relative velocity effect.  Early galaxies may be observed in the next decade, either directly via the James Webb Space Telescope or indirectly through the near-infrared background and its fluctuations.  Due to the relative velocity effect, the large-scale clustering of these galaxies should show scale-dependent bias and (if sufficient statistics are available to split the galaxy population into multiple samples) relative stochasticity between different samples of galaxies. For example, in our calculation at $z=15$, haloes of mass $10^8\\,M_\\odot$ show an increase of $\\Delta b=0.73$ over the no-velocity result of $b_0=4.6$ at $k=0.02\\,$Mpc$^{-1}$. Whether this effect will be detectable depends on the as-yet-unknown luminosity function of the highest-redshift galaxies, and whether the relative velocity effect can be separated from the scale-dependent bias produced by reionization \\cite{2006MNRAS.365.1012F, 2007MNRAS.377.1175D}.  Since reionization is believed to be driven by the formation of the first haloes massive enough to contain gas and produce stellar photoionizing radiation \\cite{2001PhR...349..125B}, the relative velocity effect will delay reionization.  However, this effect is degenerate with the mapping from halo mass to the number of massive stars formed and given the modest (factor of $\\sim 2$) effects investigated here we do not expect that the effects can be disentangled from the reionization history alone.  More interesting would be an investigation of the spatial structure of reionization and of the related high-redshift 21 cm signal, which has been investigated analytically and in simulations \\cite{2003MNRAS.341...81I, 2004ApJ...608..622Z, 2004ApJ...613....1F, 2006MNRAS.365..115F, 2006MNRAS.366..689C, 2007ApJ...654...12Z}.  The scale-dependent bias and stochasticity we find here for early haloes may have a significant effect on the structure (power spectrum, topology, and correlation with the matter density field) of early reionization bubbles.  If the {\\em pre}-reionization 21 cm signal \\cite{2004PhRvL..92u1301L} is ever observed, and cosmological information extracted, the relative velocity effect will be very important: the smallest scale fluctuations in the baryons are modified at the tens of percents level.  Indeed, since the 21 cm signal is nonlinear in the baryon density (in the limit where the hydrogen spin temperature is closely coupled to the CMB temperature, the signal is proportional to $n^2$ times the temperature-dependent collision cross section \\cite{2005ApJ...622.1356Z}), it is likely that even the large-scale fluctuations would be affected because regions with increased small-scale baryon power spectra will show more absorption. The locally anisotropic nature of the small-scale baryonic perturbations would also represent an issue for weak lensing of the 21 cm field \\cite{2004NewA....9..417P, 2006ApJ...647..719M, 2006ApJ...653..922Z, 2007MNRAS.381..447M} and/or non-Gaussianity searches \\cite{2006PhRvL..97z1301C, 2008PhRvD..77j3506C}. A full analysis of the effect on the 21 cm power spectrum and non-Gaussian statistics is deferred to future work.  Finally, the $\\Lambda$CDM cosmology predicts that early dark matter haloes in the affected range of scales (mainly $\\sim 10^4$--$10^8\\,M_\\odot$) are assimilated into larger structures.  Some of these early haloes may still be present today as dark matter substructure, which has attracted a great deal of interest since the subhalo mass function is in principle sensitive to the primordial small-scale CDM power spectrum and hence to possible deviations from ``vanilla'' CDM behavior (e.g. warm dark matter, or particles that are kinetically coupled to baryons at high redshift). Unfortunately, the overall power suppression effect we describe is probably {\\em not} detectable via substructure since the transition from an initial CDM power spectrum through the formation and survival of substructure is still not quantitatively understood (e.g. \\cite{2003ApJ...598...49Z, 2007ApJ...667..859D, 2008MNRAS.391.1685S}).  However, the power suppression effect is modulated by the relative velocity field, which comes primarily from large-scale modes in the primordial density fluctuations ($k\\sim 0.1\\,$Mpc$^{-1}$) and hence can be reconstructed from large-scale structure surveys. Therefore it would be valuable for future work to investigate whether the $v_{\\rm bc}$ effect can be detected by differential measurements that compare the substructure abundance in strong lens systems where the lens haloes have similar mass but different reconstructed ${\\bf v}_{\\rm bc}$. It is also important to mention that the suppression of the formation of the first halos, which seed present day galaxies, and the decrease in the high-k power spectrum might help alleviating the known problem of the over-abundance of substructure of dark matter halos in the $\\Lambda$CDM model (the missing satellite problem). This effect might also be important for predictions of the annihilation signal from dark matter particles. We relegate detailed investigation of these questions to the future study.  In summary, we have shown that in the post-recombination Universe, there are bulk relative motions between the baryons and dark matter that are supersonic and are coherent over scales of several comoving megaparsecs.  The combined growth of small scale structure (between the baryon Jeans length $\\lambda_{\\rm J}$ and ${\\cal M}\\lambda_{\\rm J}$, where ${\\cal M}$ is the Mach number of relative motion) is suppressed due to the baryons advecting out of the potential wells created by the dark matter.  We find at lower redshifts (e.g. $z\\sim 40$) a suppression of the power spectrum by $\\sim 10$\\% on scales of 50--500$\\,$Mpc$^{-1}$ that is highly spatially variable.  The suppression results in some reduction in the abundance of early haloes, but more importantly changes their spatial structure, leading to scale-dependent bias and stochasticity of the first haloes.  These latter effects may be large for early low-mass, high-bias haloes, e.g. we find squared correlation coefficients $\\chi=r_{\\rm hm}^2$ as small as $\\sim 0.2$ at $z=40$.  Whether this unusual clustering pattern affects the spatial morphology of reionization depends on the importance of low-mass haloes and hence is unknown at this time, although we note that future 21 cm observations combined with simulations that distribute their sources of ionizing radiation in different ways may shed some light on this issue.  Farther in the future, the power suppression effect would certainly be significant for the interpretation of any pre-reionization 21 cm signal.  In any case, our analysis highlights the importance of reconsidering ``standard'' notions of structure formation (e.g. linear bias of haloes on large scales) as we enter new physical regimes at high redshift."
    },
    "1005/1005.0625_arXiv.txt": {
        "abstract": "The rapid circularization and synchronization of the stellar components in an eccentric binary system at the onset of Roche lobe overflow is a fundamental assumption common to all binary stellar evolution and population synthesis codes, even though the validity of this assumption is questionable both theoretically and observationally. Here we calculate the evolution of the orbital elements of an eccentric binary through the direct three-body integration of a massive particle ejected through the inner Lagrangian point of the donor star at periastron.  The trajectory of this particle leads to three possible outcomes: direct accretion onto the companion star within a single orbit, self-accretion back onto the donor star within a single orbit, or a quasi-periodic orbit around the companion star, possibly leading to the formation of a disk. We calculate the secular evolution of the binary orbit in the first two cases and conclude that direct impact accretion can increase as well as decrease the orbital semi-major axis and eccentricity, while self-accretion always decreases the orbital semi-major axis and eccentricity. In cases where mass overflow contributes to circularizing the orbit, circularization can set in on timescales as short as a few per cent of the mass transfer timescale. In cases where mass overflow increases the eccentricity, the orbital evolution is governed by competition between mass overflow and tidal torques. In the absence of tidal torques, mass overflow results in direct impact can lead to substantially subsynchronously rotating donor stars. Contrary to assumptions common in the literature, direct impact accretion furthermore does not always provide a strong sink of orbital angular momentum in close mass-transferring binaries; in fact we instead find that a significant part can be returned to the orbit during the particle orbit. The formulation presented in this paper together with our previous work can be combined with stellar and binary evolution codes to generate a better picture of the evolution of eccentric, Roche lobe overflowing binary star systems. ",
        "introduction": "In close binaries, the combined effects of stellar and orbital evolution can cause a star to fill its Roche lobe and transfer mass to its companion. If the orbit is eccentric this mass transfer is expected to commence at or near the periastron of the binary orbit where the effective Roche lobes of the component stars are smallest \\citep{1998NewA....3..111L, 2007ApJ...660.1624S, 2009MNRAS.tmp..394C, 2005MNRAS.358..544R}. Until recently, studies of mass transfer in eccentric binaries mainly focused on smoothed particle hydrodynamics calculations of the mass transfer stream and much less on the effects of the mass transfer on secular binary evolution.  To the best of our knowledge, \\citet{2007ApJ...667.1170S, 2009ApJ...702.1387S} were the first to study in detail the secular evolution of the orbital semi-major axis and eccentricity of eccentric mass-transferring binaries. The authors adopted a perturbation method initially outlined by \\citet{1969Ap&SS...3..330H} and found that, depending on the binary properties at the onset of mass transfer, the orbital semi-major axis and eccentricity can either increase or decrease at a rate linearly proportional the mass transfer rate at periastron. Contrary to common assumptions, tides therefore do not always rapidly circularize binary orbits as a star evolves towards filling its Roche lobe.  \\citet{2007ApJ...667.1170S, 2009ApJ...702.1387S} calculated the orbital evolution due to mass transfer in eccentric binaries assuming mass transfer leads to the formation of an accretion disk. Tidal interactions between the disk and the donor star then allow at least some fraction of the angular momentum carried by the transferred matter to be returned to the orbit. However, if the size of the accretor is a considerable fraction of the size of the orbit, matter lost by the donor star through the inner Lagrangian point may impact the accretor's surface directly rather than form an accretion disk.  In this paper, we complement the work of \\citet[][ hereafter Paper~I]{2007ApJ...667.1170S} and \\citet{2009ApJ...702.1387S} by investigating the orbital evolution due to mass overflow in eccentric binaries in which no disk forms around the accretor. For this purpose, we calculate the trajectories of mass elements escaping through the inner Lagrangian point of the donor star in the ballistic limit, and distinguish between three possible cases: (1) the ejected matter falls back onto the donor star, (2) the ejected matter impacts the surface of the accretor, and (3) the ejected matter orbits the accretor for at least one orbital revolution of the binary.  In the first two cases, the evolution of the orbital elements is obtained directly from the ballistic trajectory calculations. In the third case, the evolution of the orbital elements depends on the interactions between mass elements ejected during successive periastron passages of the donor star. Since such interactions are expected to violate our assumption of particles moving in the ballistic limit, we do not consider them here.  The present paper is organized as follows.  In \\S\\,2, we outline the basic assumptions and methodology used to study ballistic particle trajectories and the effects of mass overflow on the orbital evolution of eccentric binaries. In \\S\\,3, we examine the fate of mass lost by the donor star as a function of the initial binary properties and identify ranges of initial parameters leading to fallback, direct impact accretion, and possible disk formation. Evolutionary sequences and timescales of orbital evolution due to mass overflow in eccentric binaries are presented in \\S\\,4. The final section, \\S\\,5, is devoted to concluding remarks. ",
        "conclusions": "In this paper, we extended our previous work on mass transfer in eccentric binaries by using ballistic trajectory calculations to determine the parameter space for different outcomes of mass overflow. Assuming mass overflow takes place through Roche lobe overflow at the periastron of the binary orbit, we explored a broad parameter space to determine the conditions under which mass overflow leads to direct impact accretion onto the donor's companion or to fallback onto the donor star (``self-accretion''). The results presented in this paper along with those presented in \\citet{2007ApJ...660.1624S, 2007ApJ...667.1170S, 2009ApJ...702.1387S}, provide a self-consistent picture of the orbital evolution of eccentric effective Roche lobe overflowing binary star systems as a function of their initial orbital parameters. Until now, such a picture has been lacking both in binary stellar evolution and binary population synthesis codes.  We find that systems with a large initial eccentricity ($e_i \\gtrsim 0.2$) or a donor initially rotating near synchronicity ($f_{1,i} \\gtrsim 0.8$) are always expected to undergo self-accretion. Direct impact accretion or disk formation is expected to occur mainly for systems with low eccentricities, low mass ratios, and substantially subsynchronously rotating donor stars.  Furthermore, self-accretion is found to always decrease the orbital eccentricity, and can do so over timescales ranging from less than a Myr to more than a Gyr depending on the initial binary parameters and the chosen mass overflow rate.  To illustrate the applicability of the presented formalism, we calculated the orbital evolution due to mass overflow for eccentric binaries consisting of an evolved main sequence donor and a zero-age main sequence accretor as a function of the initial binary parameters. For binaries undergoing direct impact accretion, mass transfer can increase as well decrease the initial orbital semi-major axis and eccentricity, while for binaries undergoing self-accretion, mass overflow always decreases both the initial orbital semi-major axis and eccentricity. For a mass overflow rate of $\\dot{M}_0=-10^{-9}\\,M_\\odot\\, {\\rm yr}^{-1}$, the time scales of orbital evolution can range from less than 1\\,Gyr to more than a Hubble time.  In this exploratory study, both the donor star and the accretor were treated as rigid spheres. The results presented therefore do not account for orbital evolution due to tidal torques or for the radial response of the donor star to mass loss. Our calculations nevertheless clearly illustrate that mass overflow leading to direct impact accretion or self-accretion can both enhance or counteract the effects of tides in eccentric effective Roche lobe overflowing binaries. A similar conclusion was reached by \\citet{2007ApJ...667.1170S, 2009ApJ...702.1387S} for binaries in which mass transfer leads to the formation of an accretion disk. As one of the next steps in our investigation, we will implement more realistic stellar models in our ballistic particle trajectory code and examine the effects of tides and stellar evolution on mass transfer in eccentric binaries.  While not discussed in detail here, we also find that direct impact accretion does not necessarily deposit all of the accreted matter's angular momentum into the spin of the accretor. Instead, some of this angular momentum, which originated from both spin and orbital angular momentum of the donor star, is deposited into the accretor's spin while the remaining part is returned to the orbit. This is in stark contrast to commonly adopted assumptions in stability studies of mass transfer in ultra-compact binaries \\citep[e.g.,][]{1988ApJ...332..193V, 2004MNRAS.350..113M}. The possibility of avoiding strong orbital angular momentum loss during direct impact accretion is particularly relevant in predicting expected gravitational wave detection rates from mass-transferring double white dwarfs (AM\\,CVn stars) by the Laser Interferometer Space Antenna (LISA)."
    },
    "1005/1005.2620_arXiv.txt": {
        "abstract": "Cosmogenic neutrinos originate from photo-hadronic interactions of cosmic ray protons with the cosmic microwave background (CMB).  The neutrino production rate can be constrained through the accompanying electrons, positrons and $\\gamma$-rays that quickly cascade on the CMB and intergalactic magnetic fields to lower energies and generate a $\\gamma$-ray background in the GeV-TeV region.  Bethe-Heitler pair production by protons also contributes to the cascade and can tighten the neutrino constraints in models where extragalactic cosmic rays begin to dominate over the galactic component at a relatively low ``crossover'' energy.  We investigate this issue in the light of the recent Fermi-LAT measurements of the diffuse extragalactic $\\gamma$-ray background and illustrate by a fit to the HiRes spectrum how the prediction of the cosmogenic neutrino flux in all-proton models varies with the crossover energy.  The neutrino flux is required to be smaller when the $\\gamma$-ray bound is applied, nevertheless such models are still consistent with HiRes and Fermi-LAT if one properly takes into account the energy uncertainty of cosmic ray measurements. The presently allowed flux is within reach of the IceCube neutrino telescope and other dedicated radio experiments. ",
        "introduction": "Soon after the discovery of the cosmic microwave background (CMB)~\\cite{Penzias:1965wn}, it was realized that interactions of extragalactic ultrahigh energy (UHE) cosmic rays (CRs) on the relic photons would suppress the cosmic ray flux at energies $\\gtrsim 5 \\times 10^{10}$~GeV, the so-called ``GZK cutoff''~\\cite{Greisen:1966jv,Zatsepin:1966jv}. It was pointed out subsequently~\\cite{Beresinsky:1969qj} that the GZK interaction also generates a ``cosmogenic flux'' of neutrinos, through the decay of secondary charged pions. Forty years later, the predicted suppression of the UHE CR flux was indeed observed by the HiRes~\\cite{Abbasi:2007sv} and Auger~\\cite{Abraham:2008ru} experiments. However, the cosmogenic flux of neutrinos has yet to be detected.  The GZK reaction chain generating cosmogenic neutrinos is well known~\\cite{Stecker:1978ah}. The intermediate state of the reaction $p \\gamma_{\\rm CMB} \\to n \\pi^+/p\\pi^0$ is dominated by the $\\Delta^+$ resonance, because the neutron decay length is smaller than the nucleon mean free path on the CMB. Resonant $p\\gamma$ interactions produce twice as many neutral pions as charged pions.  Direct pion production via virtual meson exchange contributes only about 20\\% to the total cross-section, but is almost exclusively into $\\pi^+$. Hence, $p\\gamma$ interactions produce roughly equal number of $\\pi^+$ and $\\pi^0$. Gamma-rays, produced via $\\pi^0$ decay, subsequently cascade electromagnetically on intergalactic radiation fields through $e^+ e^-$ pair production followed by inverse Compton scattering. The net result is a pile up of $\\gamma$-rays at GeV-TeV energies, just below the threshold for further pair production on the diffuse optical background. Meanwhile each $\\pi^+$ decays to 3 neutrinos and a positron; the $e^+$ readily loses its energy through inverse Compton scattering on the diffuse radio background or through synchrotron radiation in intergalactic magnetic fields. The neutrinos carry away about 3/4 of the $\\pi^+$ energy, therefore the energy in cosmogenic neutrinos is about 3/4 of that produced in $\\gamma$-rays.  The normalization of the neutrino flux depends critically on the cosmological evolution of the CR sources and on their proton injection spectra~\\cite{Yoshida:pt}. It also depends on the assumed spatial distribution of sources; for example, local sources in the Virgo cluster~\\cite{Hill:1985mk}, would dominate the high energy tail of the proton spectrum.  Another source of uncertainty is the energy at which there is a transition from Galactic to extragalactic CRs as inferred from a change in the spectral slope.  The ``ankle'' at $\\sim 3 \\times10^9$~GeV seems to be a natural candidate for this transition~\\cite{Linsley:1963bk,Hill:1983mk,Wibig:2004ye}, but a lower energy crossover at the ``second knee'' at $\\sim 5 \\times10^8$~GeV has also been advocated~\\cite{Berezinsky:2002nc,Fodor:2003ph}.  A fourth source of uncertainty is the chemical composition of the parent CRs -- if these are heavy nuclei rather than protons, then the neutrino flux is reduced~\\cite{Hooper:2004jc}.  The most up-to-date calculation~\\cite{Anchordoqui:2007fi} of the cosmogenic neutrino flux combines a double-fit analysis of the energy~\\cite{Abraham:2010mj} and elongation rate~\\cite{Abraham:2010yv} measurements to constrain the spectrum and chemical composition of UHE CRs at their sources. Injection models with a wide range of chemical compositions are found to be consistent with observations. The data is consistent with a proton-dominated spectrum with a small admixture of heavy nuclei, in which case the cosmogenic neutrino flux is rather similar to the all-proton model. In this case, kilometer-scale neutrino telescopes are expected to observe of ${\\cal O}(1)$ cosmogenic neutrino event per year. In contrast, an intermediate to heavy nuclear composition beyond the ankle, as indicated by the elongation rate vs.~energy~\\cite{Abraham:2010yv}, can lead to a considerable suppression (up to two orders of magnitude) of the cosmogenic neutrino flux in comparison to the all-proton case. However, little is known about the chemical composition from just below to beyond the GZK cutoff, where the most significant contribution to cosmogenic neutrinos form  UHE CR protons is expected. It is interesting to note that uncertainties in the extrapolation of the proton-air interaction -- cross-section, elasticity and multiplicity of secondaries -- from accelerator measurements to the high energies characteristic for air showers are large enough to undermine any definite conclusion on the chemical composition~\\cite{Wibig:2008ji,Ulrich:2009yq}.  In this work we study the constraint set by the diffuse $\\gamma$-ray background on all-proton models of extragalactic CRs.  We parametrize our ignorance of the crossover energy -- which marks the transition between the galactic and extragalactic components -- as a variable low energy cutoff in the proton injection rate. By fitting only to CR data above the crossover energy, taken to be between $10^{17.5}$~eV and $10^{19}$~eV, we determine the statistically preferred values of the spectral index $\\gamma$ and cosmic source density evolution index $n$ by a goodness-of-fit (GOF) test of the HiRes data, taking into account the energy resolution of about 25\\%. For each model we check that the total energy density of the EM cascade is below a critical value inferred from the recent measurement of the extragalactic $\\gamma$-ray background by the Fermi-LAT Collaboration~\\cite{Abdo:2010nz}.  We find that the allowed range of the cosmogenic neutrino flux increases with the crossover energy and can be up to an order of magnitude {\\em larger} than the values presented in a recent study~\\cite{Berezinsky:2010xa}.  This paper is organized as follows. We begin in \\S~\\ref{sec:I} with a discussion of the extragalactic proton fluxe and the corresponding energy density of cascade $\\gamma$-rays. We present our statistical method in \\S~\\ref{sec:II} and discuss our results in \\S~\\ref{sec:III} ",
        "conclusions": "\\label{sec:III}  The cosmogenic neutrino fluxes that we have shown in Fig.~\\ref{fig:comparison} are compared to present upper limits on the diffuse neutrino flux in Fig.~\\ref{fig:nulimits}. As before, the solid green line shows the neutrino flux (summed over flavours) corresponding to the best fit of the proton spectra and the dashed green line indicate the range of neutrino fluxes within the 99\\%~C.L. For all crossover energies considered, the range of models at the 99\\% C.L.~is consistent with existing neutrino limits. For illustration, the thin dotted line shows the larger range of neutrino fluxes at the 99\\% C.L.~corresponding to a fit {\\em without} the Fermi LAT constraint ({\\it cf.}~the black contours in the left panel of Fig.~\\ref{fig:gof}). It is apparent that this indirect bound from GeV-TeV $\\gamma$-rays does reduce the number of possible models significantly.  At this point it is worth stressing that the Fermi-LAT spectrum used in this analysis is not the result of a direct observation but is derived by a foreground subtraction scheme. The extra-galactic $\\gamma$-ray background inferred by EGRET~\\cite{Sreekumar:1997un} shows a significantly larger intensity and a harder spectral index. A possible source of the differences could be due to the different diffuse galactic emission (DGE) models used in the analysis. As pointed out in~\\cite{Abdo:2010nz} a re-analysis of the EGRET data with an updated DGE model~\\cite{Strong:2004ry} is comparable with the intensity observed with Fermi-LAT. It is beyond the scope of this paper to address these systematic uncertainties. The {\\it maximal} effect of a larger $\\gamma$-ray background intensity is indicated by the extended parameter regions shown in Fig.~\\ref{fig:gof} which are derived {\\em without} the Fermi-LAT constraint together with the corresponding range of neutrino fluxes in Fig.~\\ref{fig:nulimits}.  The overall range of neutrino fluxes increases along with the crossover energy - not only in magnitude, which is expected already due to the reduced set of CR data used in the GOF test, but also to significantly larger neutrino fluxes. Also the cosmogenic neutrino flux of the best-fit models increases by over a factor of two in the peak region ($\\sim 10^9$~GeV). This confirms our earlier suspicion that an increasing value of the crossover energy allows a larger contribution of cosmogenic neutrinos relative to the $\\gamma$-rays and hence larger neutrino fluxes.  Figure~\\ref{fig:nulimits} also shows the estimated sensitivity of IceCube~\\cite{Ahrens:2003ix} (5$\\sigma$) to neutrino fluxes in the $10^6$-$10^8$~GeV (solid)~\\cite{Achterberg:2006md} and the $10^8$-$10^{10}$~GeV (dotted)~\\cite{Halzen:2006ic} energy range after one year of observation. IceCube located at the South Pole is presently the largest neutrino telescope. On completion in early 2011 it will consist of a km${}^3$-scale detector of transparent glacial ice, that is constantly monitored for \\v{C}erenkov light emission of secondary charged particles from high energy neutrino interactions. It is apparent from Fig.~\\ref{fig:nulimits} that IceCube's sensitivity after one year is already sufficient to probe cosmogenic neutrino fluxes from an all-proton spectrum of extra-galactic cosmic rays. If the crossover energy exceeds $10^{18.5}$~eV, the best-fit model of the HiRes data is within reach of IceCube.  In summary we find that while the expected range of cosmogenic neutrino fluxes in all-proton models is indeed reduced due to the constraint from the diffuse $\\gamma$-ray background measurements, neutrino fluxes compatible at 99\\% C.L with the HiRes and Fermi-LAT results can be larger than those presented in Ref.\\cite{Berezinsky:2010xa} by up to factor of $\\sim 30$ for the same values of $E_{\\rm max}=10^{21}$ eV and $z_{\\rm max}=2$. In particular the allowed cosmogenic flux is still within reach of neutrino observatories like IceCube.  Furthermore, our results are obtained with the simple parametrization of the source spectral emission rate in Eq.~\\eqref{eq:injection}; larger neutrino fluxes might be allowed with a more general spectrum than a simple power-law.  One can also turn this argument around and use observation or non-observation of cosmogenic neutrinos in the near future to provide additional constraints on the composition of cosmic rays~\\cite{Ahlers:2009rf}. We have assumed here an all-proton composition for extra-galactic cosmic rays. However, as mentioned already, the chemical composition of UHE CRs is rather uncertain and may well be dominated by heavy nuclei. In this case the limits on diffuse neutrino fluxes can still serve as a probe of the possible proton fraction in cosmic rays~\\cite{Ahlers:2009rf,Anchordoqui:2009nf} and the limits from diffuse $\\gamma$-rays serve as an additional probe~\\cite{Ahlersetal}."
    },
    "1005/1005.0555_arXiv.txt": {
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        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.0249_arXiv.txt": {
        "abstract": "{ During its 6 years of operation, \\emph{INTEGRAL}/ISGRI has detected more than 500 sources. Many of these sources are variable. Taking into account that nearly half of \\emph{INTEGRAL}/ISGRI sources are new and many of them remain unidentified, the variability properties of the sources can provide additional constraints to help us to classify and identify the unknown sources. } { To study the variability properties of the sources detected by \\emph{INTEGRAL}/ISGRI, we develop a method to quantify the variability of a source. We describe here our techniques and compile a catalog of the sources that fit our criteria of variability. } { We use the natural time binning of \\emph{INTEGRAL} observations called Science Window ($\\approx 2000$ seconds) and test the hypothesis that the detected sources are constant using a $\\chi^2$ all-sky map in three energy bands (20-40, 40-100, 100-200 keV). We calculate an intrinsic variance of the flux in individual pixels and use it to define the fractional variability of a source. The method is sensitive to the source variability on timescales of one Science Window and higher. We concentrate only on the sources that were already reported to be detected by \\emph{INTEGRAL}. } { We present a catalog of 202 sources found to be significantly variable. For the catalog sources, we give the measure of variability and fluxes with corresponding errors in the 20-40, 40-100 and 100-200 keV energy bands, and we present some statistics about the population of variable sources. The description of the physical properties of the variable sources will be given in a forthcoming paper. } {} ",
        "introduction": "\\emph{INTEGRAL} (INTErnational Gamma Ray Astrophysics Laboratory, \\citet{Winetal03}) was launched in 2002 and since then has performed high-quality observations in the energy band from 3 keV up to $\\sim 10$ MeV. The \\emph{INTEGRAL} payload consists of two main soft gamma-ray instruments (the imager IBIS \\citep{Ubeetal03}, and the spectrometer SPI \\citep{Vedetal03}) and two monitors (in X-rays JEM-X \\citep{Lunetal03}, and in optical OMC \\citep{Masetal03}). The wide field-of-view of the imager IBIS provides an ideal opportunity to survey the sky in hard X-rays.\\\\  During its first 6 years in orbit, \\emph{INTEGRAL} has covered nearly the whole sky. The observational data have been mainly used to study the soft gamma-ray emission from the Galactic plane (GP) \\citep{Bouetal05, Krietal06} through the Galactic plane scans and the Galactic centre (GC) \\citet{Beletal04, Revetal04, Beletal06, Krietal06} through the Galactic centre deep exposure programme. A number of papers have already presented general surveys \\citep{Bazetal06, Biretal06} of the sky as well as of specific regions \\citep{Gotetal05, Haretal06, Moletal04} and population types \\citep{Baretal06, Basetal06, Sazetal07, Becetal06, Becetal09}.\\\\  The majority of the classified sources detected by \\emph{INTEGRAL} are either low and high mass X-ray binaries (LMXBs and HMXBs) or AGNs \\citep{Bodetal07}. However, a significant fraction of the detected sources remain unidentified. A special approach to population classification is required for the GC region to resolve the population types because of the high density of sources. Fortunately, the physics of the sources may help us to unveil their type. Indeed, the bulk of the \\emph{INTEGRAL} sources are accreting systems that are expected to be intrinsically variable on multiple timescales depending on the source type and the nature of the variability. For instance, X-ray binaries (XRBs) may exhibit variability on timescales that range from milliseconds (supporting the idea that emission originates close to the compact object in the inner accretion radius) to hours and days, indicating that the variability can originate throughout the accretion flow at multiple radii and propagate inwards to modulate the central X-ray emission \\citep{AreUtt05}. This idea is supported by the known correlation between millisecond/second and hour/day scale variability in XRBs \\citep{Utt03}. LMXBs may exhibit flaring behavior with an increase in both emission intensity and hardness over a period of a few hundred to a few thousand seconds. X-ray bursts with rise times of a few seconds and decay times of hunderds of seconds or even several hours \\citep{Baretal02} are also common to these objects. On the other hand, HMXBs are known to exhibit variability on timescales ranging from a fraction of a day up to several days, generated by the clumpiness of the stellar wind acreting onto the compact object \\citep{Ducetal09}. Hour-long outbursts caused by variable accretion rates are observed in supergiant fast X-ray transients, a sub-class of HMXBs discovered by INTEGRAL \\citep{Rometal2009}. Owing to their larger size, AGNs of different types exhibit day-to-month(s) variability depending on the black hole mass \\citep{IshCou09}. Gamma-ray loud blazars have variability timescales in the range from $10^{1.6}$ to $10^{5.6}$~s \\citep{LiaLiu03}. Therefore, a list of \\emph{INTEGRAL} sources with quantitative measurements of their variability would be an important help to classifying the unidentified sources and more detailed studies of their physics.\\\\  The variability of \\emph{INTEGRAL} sources was addressed in the latest 4th IBIS/ISGRI survey catalog paper \\citep{Biretal09} when the authors performed the so-called \\textit{bursticity} analysis intended to facilitate the detection of variable sources.\\\\  Here we present a catalog of \\emph{INTEGRAL} variable sources identified in a large fraction of the archival public data. In addition to standard maps produced by the standard data analysis software, we compiled a $\\chi^2$ all-sky map and applied the newly developed method to measure the fractional variability of the sources detected by the IBIS/ISGRI instrument onboard \\emph{INTEGRAL}. The method is sensitive to variability on timescales longer than those of single ScW exposures ($\\approx 2000$ seconds), i.e., to variability on timescales of hour(s)-day(s)-month(s). The catalog is compiled from the sources detected in the variability map. In addition, we implemented an online service providing the community with all-sky maps in the 20-40, 40-100, and 100-200 keV energy bands generated during the course of this research.\\\\  In the following, we describe the data selection procedure and the implemented data analysis pipeline (Sect.~\\ref{sec:datana}). In Sect.~\\ref{sec:method}, we outline our systematic approach to the detection of variability in \\emph{INTEGRAL} sources and describe our detection procedure in Sect.~\\ref{sec:detect}. We compile the variability catalog in Sect.~\\ref{sec:catvar}. In Sect.~\\ref{sec:skyview}, we briefly describe the implemented all-sky map online service. We make some concluding remarks in Sect.~\\ref{sec:conclu}. ",
        "conclusions": ""
    },
    "1005/1005.2489_arXiv.txt": {
        "abstract": "{} {Our aim is to confirm the interstellar detection of cyanic acid, HOCN, in the Galactic center clouds. It has previously been tentatively detected only in Sgr B2(OH).} {We used a complete line survey of the hot cores Sgr B2(N) and (M) in the 3 mm range, complemented by additional observations carried out with the IRAM 30~m telescope at selected frequencies in the 2 mm band and towards four additional positions in the Sgr B2 cloud complex in the 2 and 3 mm bands. The spectral survey was analysed in the local thermodynamical equilibrium approximation (LTE) by modeling the emission of all identified molecules simultaneously. This allowed us to distinguish weak features of HOCN from the rich line spectrum observed in Sgr B2(N) and (M). Lines of the more stable (by 1.1~eV) isomer isocyanic acid, HNCO, in these sources, as well as those of HOCN and HNCO towards the other positions, were analysed in the LTE approximation as well.} { Four transitions of HOCN were detected in a quiescent molecular cloud in the Galactic center at a position offset in (R.A., decl.) by (20$^{\\prime\\prime}$,100$^{\\prime\\prime}$) from the hot core source Sgr B2(M), confirming its previous tentative interstellar detection. Up to four transitions were detected toward five other positions in the Sgr B2 complex, including the hot cores Sgr B2(M), (S), and (N). A fairly constant abundance ratio of $\\sim 0.3-0.8$~\\% for HOCN relative to HNCO was derived for the extended gas components, suggesting a common formation process of these isomers. } {} ",
        "introduction": "Ever since its first detection towards Sgr B2 (OH) \\citep{SB1972}, isocyanic acid (HNCO) has been found with high abundances in a multitude of galactic sources, such as high-mass star-forming regions \\citep{JAB1984,ZHM2000,BJD2007}, Galactic center molecular clouds \\citep{DHW1997,MRM2008}, as well as translucent and dark clouds \\citep {TTH1999}. HNCO has also been detected in several external galaxies \\citep{NHJ1991,WHC2004,MMM2009}, and been suggested as a tracer for nuclear starburst activity \\citep{MRM2008,MMM2009}. Two energetically less stable isomers of HNCO have recently been detected in the interstellar medium: cyanic acid (HOCN) tentatively towards Sgr B2 (OH) following its laboratory identification \\citep{BGM2009b}, and fulminic acid (HCNO) towards three starless dark clouds and one low-mass star-forming region \\citep{MCT2009} with abundance ratios relative to HNCO of 0.5\\% and $1-5$\\%, respectively. The relative abundances of members of an isomeric family reflect their different formation and destruction processes, and observations of isomers of  small systems like HCN/HNC  \\citep[e.g., ][]{SWP1992} as well as complex organic species \\citep{HLJ2000,TA2001} have been previously used to investigate their gas-phase and/or grain-surface interstellar origin.  The Giant Molecular Cloud Sgr B2 at a projected distance of $\\sim$90~pc from the Galactic center, where HOCN had been tentatively detected, harbors three massive star-forming regions associated with the most massive (sub)mm continuum sources in this region: Sgr B2(N), (M), and (S). These hot cores are aligned in a north-south direction in the dense central $5-10$ pc of Sgr B2 and have been shown to be extremely rich in organic molecules by extensive molecular line surveys \\citep[e.g., ][]{CLT1986,Turner1991,NBH2000,BMC2008a,BMC2008b,BGM2009a}.  Whereas most molecules have peak abundances in the hot cores, HNCO emission was found to be extended with an expanding ringlike morphology and a peak at a position 2' north (2'N) of Sgr B2(M), where only a few other molecules, like HCO$_2^+$ and CH$_3$OH, have enhanced abundances \\citep{MHH1998,MI2006,JBC2008}. The HNCO enhancement at this particular position has been explained by a cloud-cloud collision, resulting in shocks that release grain mantle material which is further processed via gas-phase reactions \\citep{MI2006}. Similarly, large-scale, low-velocity shocks are proposed as responsible for the large abundances of complex organic molecules, including HNCO, in the quiescent molecular clouds throughout the whole Galactic-center region \\citep{RMR2006,RMM2008}.  In this paper we present additional astronomical observations of HOCN toward several positions in the Sgr B2 complex with the IRAM 30m telescope. These observations confirm the tentative identification by \\cite{BGM2009b} based on previously published line surveys of Sgr B2(OH). We find evidence of HOCN towards six positions, including the hot cores Sgr B2(M) and (N). The interstellar detection of HOCN is unequivocally confirmed by the observation of four unblended transitions toward a position very close to the 2'N cloud, at an offset in (R.A., decl.) by (20$^{\\prime\\prime}$,100$^{\\prime\\prime}$) from Sgr B2(M). We compare the relative abundances of HOCN to the most stable isomer HNCO. ",
        "conclusions": "HOCN has previously been tentatively detected \\citep{BGM2009b} based on published spectral line survey data of Sgr B2(OH) obtained with the BTL 7~m and NRAO 11~m telescopes \\citep{CLT1986,Tur1989}. The current observations confirm and improve the detection of HOCN in several important ways: i) Four unblended lines of HOCN were detected in the extended gas at an offset position (20$^{\\prime\\prime}$,100$^{\\prime\\prime}$) from the hot core Sgr B2(M) close to the maximum of HNCO emission. This region does not suffer from the severe line confusion affecting the hot core regions, making incidental blendings and false assignments less likely. ii) The comparison of the observed spectrum with a synthetic spectrum -- based on the simultaneous analysis of our complete 3~mm survey data -- allows disentangling contributions of known, mostly complex, molecular species from the weak features associated with HOCN in the hot core regions Sgr B2(N) and (M). iii) The beamsize of the IRAM 30~m telescope is 2.7-times smaller than that of the NRAO 11~m telescope used by \\cite{Tur1989}. This higher spatial resolution allows separation of different components in a complex source like Sgr B2. iv) The observed line intensity ratios are consistent with those predicted in the LTE approximation. v) It should also be noted that all lines of HOCN with an expected intensity over the noise level have either been detected or are blended  in the covered frequency range towards the six selected sources; i. e., no expected lines are missing.  HNCO is ubiquitous throughout the Central Molecular Zone (CMZ) of our Galaxy \\citep{Menten2004}, stretching from longitudes of $-0.2$ to 1.8 degrees over a $\\sim$300~pc long band along the Galactic plane \\citep{DHW1997}. In Sgr B2 -- the most prominent molecular emission region within the CMZ -- the molecule exhibits large-scale abundance variations peaking not at the hot core positions -- where most other molecules show strong emission -- but rather has an extended ring-like structure, investigated in detail by \\cite{MI2006} and \\cite{JBC2008}. The enhancement may result from shock evaporation of icy grain mantles caused by the collision between the principal cloud and the ring, which appears to be expanding. Detection of cyanic acid toward several chemically and physically distinct regions in Sgr B2 with a fairly constant abundance ratio relative to the most stable isomer indicates an effective formation process closely related to that of HNCO. On the basis of this limited sample, HOCN may be as widespread as HNCO.  Current astrochemical gas-grain model calculations efficiently produce HNCO on grain surfaces by hydrogenation of accreted OCN \\citep{HH1993,GWH2008}. The observed fairly constant abundance ratio of HOCN relative to HNCO  in our sample of sources, which is far from thermodynamical equilibrium considering the low temperatures prevalent in the interstellar medium\\footnote{HOCN has been calculated to be around 1.1~eV higher in energy than HNCO \\citep{ML2008}, even at 200~K this results in an equilibrium ratio for HNCO/HOCN of $6\\times 10^{27}$.}, suggests that HOCN is formed along with HNCO on the grain mantles, and then subsequently released through shocks or during warm-up. Motivated by the recent detections of HCNO \\citep{MCT2009} and HOCN \\citep[][and this work]{BGM2009b}, E. Herbst and coworkers are currently extending the Ohio State gas-grain network to account for the formation and destruction of HNCO and its higher lying isomers in different types of sources (Quan {\\it et al.}, in preparation). Their hot core and warm envelope models can account reasonably well, i.e. with an order-of-magnitude agreement,  for the relative and absolute abundances of HNCO and HOCN observed here in the respective Sgr B2 components.  Discrepancies might come from shock chemistry (see above), which has so far not been implemented in the model calculations. From the present observations, it cannot be conclusively decided whether grain-surface chemistry plays a dominant role in the synthesis of CHNO isomers. Pure gas-phase chemical models, based on the steady-state approach used by \\cite{MCT2009}, and extended by reactions involving protonated forms of NCO, CNO, and of the four stable CHNO isomers, predict HNCO to HOCN abundance ratios matching the observed values stated in Table \\ref{table_astropar} for the Galactic center clouds \\citep{MBC2010}  The abundances of HNCO and its isomers are governed by a complex network of surface and gas-phase formation and destruction reactions. Further observations of HOCN toward additional regions of the interstellar medium rich in HNCO such as other high-mass star-forming regions, young stellar objects (YSOs), and dense cloud cores -- and the observation of HOCN in sources where HCNO has been detected and vice versa -- will help to further constrain the interstellar chemistry of this isomeric system. These studies might also be extended to related isomeric families. The higher energy isomers of the insterstellar molecule HNCS, which is isoelectronic to HNCO, have recently been detected in the laboratory \\citep[HSCN;][]{BYG2010}, (HCNS and HSNC; M. C. McCarthy, in preparation), and HSCN has now also been detected in Sgr B2 \\citep{HZB2009} ."
    },
    "1005/1005.0390_arXiv.txt": {
        "abstract": "In this work, decision tree learning algorithms and fuzzy inferencing systems are applied for galaxy morphology classification. In particular, the CART, the C4.5, the Random Forest and fuzzy logic algorithms are studied and reliable classifiers are developed to distinguish between spiral galaxies, elliptical galaxies or star/unknown galactic objects. Morphology information for the training and testing datasets is obtained from the Galaxy Zoo project while the corresponding photometric and spectra parameters are downloaded from the SDSS DR7 catalogue. ",
        "introduction": "\\label{IntroductionSection}  The exponential rise in the amount of available astronomical data has and will continue to create a digital world in which extracting new and useful information from archives is and will continue to be a major endeavour in itself. For instance, the seventh data release of the Sloan Digital Sky Survey (SDSS DR7), catalogues five-band photometry for 357 million distinct objects covering an area of 11,663 deg\\textsuperscript{2}, over $1.6$ million spectra for 930,000 galaxies, 120,000 quasars, and 460,000 stars over 9380 deg\\textsuperscript{2} \\citep{abazajian10}. Future wide-field imaging surveys will capture invaluable images for hundreds of millions of objects even those with very faint magnitudes.  Most of our current knowledge on galaxy classification is based on the pioneering work of several dedicated observers who visually catalogued thousands of galaxies. For instance, in \\cite{fukugita07}, 2253 objects from the SDSS DR3 were classified into a Hubble Type catalogue by three people independently.  Then, a final classification was obtained from the mean. Classifying all objects captured in very large datasets produced by digital sky surveys is beyond the capacity of a small number of individuals. This therefore calls for new approaches. The challenge here is to design intelligent algorithms which will reproduce classification to the same degree as that made by human experts. Automated methods which make use of artificial neural networks have already been proposed by \\cite{ball04} and more recently by \\cite{banerji09}. \\cite{andraeXX} also propose the use of shapelet decomposition to model the natural galaxy morphologies.  In this study, the advantages obtained by performing decision tree learning and fuzzy inferencing for galaxy morphology classification are investigated. In particular, results from the CART, the C4.5 and the Random Forest tree building algorithms are compared. The outputs obtained after testing a generated fuzzy inference system through subtractive clustering, are also presented. \\cite{ball04} adopted decision tree approaches for star/galaxy classification. \\cite{calleja04} also attempted to construct classifiers from parameters outputted from processed images of galaxies. However, such work was only constructed on a limited number of attributes and samples.  The aim here is to develop reliable models to distinguish between spiral galaxies, elliptical galaxies or star/unknown galactic objects. This through photometric as well as spectra data. Classified training and testing samples are obtained from the Galaxy Zoo project while the corresponding parameters are downloaded from the SDSS DR7 catalogue. As discussed by \\cite{banerji09}, this and similar studies present us the unique opportunity to compare human classifications to those obtained through automated machine learning algorithms. Should the automated techniques prove to be as successful as the human techniques in separating astronomical objects into different morphological classes, considerable time and effort will be saved in future surveys whilst also ensuring uniformity in the classifications.  In the following Section, the Galaxy Zoo Catalogue is described. Section \\ref{decisionTreesSection} explores the various decision tree algorithms while Section \\ref{fuzzyInferenceSystemsSection} attempts to introduce fuzzy inference systems. In Section \\ref{inputParametersSection}, the SDSS photometric and spectra parameters used for classification are discussed. The results obtained and the overall conclusion are then given in Section \\ref{resultsSection} and Section \\ref{conclusionSection} respectively. ",
        "conclusions": "\\label{conclusionSection}  In this study, accuracies for different galaxy morphology classification models developed through various machine learning techniques were obtained and analyzed. Results from the CART, the C4.5 and the Random Forest decision tree algorithms as well as the output from Fuzzy Inference Systems, were compared. The advantages gained by performing computations on different photometric parameters and on spectra attributes, were also investigated and put forward. Figure \\ref{resultsGeneralSummary} serves as a good summary of which data and algorithms were used as well as the overall accuracies obtain. In all cases, the Random Forest gave the highest percentages especially when 50 trees were used.  \\begin{figure} \\centering \\includegraphics[width=70mm]{results_generalSummary.png} \\caption {Data, algorithms and results} \\label{resultsGeneralSummary} \\end{figure}  All of the tested algorithms took only a few minutes to run on a normal personal computer. Although the presented results are for Set 3, experiments on Set 2 that stored more samples were also carried out. Accuracy percentages very close to the ones published, were obtained.  In most cases, when processing photometric parameters, the adaptive shape measure (\\texttt{mRrCc}) parameter was chosen as the root of the tree. First level nodes included the concentration (\\texttt{petroR90}/\\texttt{petroR50}) and the \\texttt{dered\\_g}-\\texttt{dered\\_r} parameters. For spectra data, the Ha wave line was determined to provide the highest information gain while the Hb and the K lines were chosen as first level nodes.  Figure \\ref{incorrectClassification} shows samples of incorrectly classified galaxies by the fuzzy inference system. Although this is not the most accurate technique described, the incorrectly classified spiral and elliptical samples are very faint in magnitude. Moreover, all incorrectly classified unknown objects have bright sources in the vicinity and this could have had an effect on the calculated parameters by the SDSS photometric pipeline.  \\begin{figure} \\centering \\includegraphics[width=84mm]{incorrect_classification.png} \\caption {Samples of spiral (top), elliptical (middle) and unknown (bottom) galaxies that were incorrectly classified by the fuzzy inference system} \\label{incorrectClassification} \\end{figure}"
    },
    "1005/1005.0673_arXiv.txt": {
        "abstract": "This paper presents an investigation into the gravitational microlensing of quasars by stars and stellar remnants in the Milky Way. We present predictions for the  all-sky microlensing optical depth, time-scale distributions and event rates for future large-area sky surveys. As expected, the total event rate increases rapidly with increasing magnitude limit, reflecting the fact that the number density of quasars is a steep function of magnitude. Surveys such as Pan-STARRS and LSST should be able to detect more than ten events per year, with typical event durations of around one month. Since microlensing of quasar sources suffers from fewer degeneracies than lensing of Milky Way sources, they could be used as a powerful tool for recovering the mass of the lensing object in a robust, often model-independent, manner. As a consequence, for a subset of these events it will be possible to directly `weigh' the star (or stellar remnant) that is causing the lensing signal, either through higher order microlensing effects and/or high-precision astrometric observations of the lens star (using, for example, Gaia or SIM-lite). This means that such events could play a crucial role in stellar astronomy. Given the current operational timelines for Pan-STARRS and LSST, by the end of the decade they could potentially detect up to 100 events. Although this is still too few events to place detailed constraints on Galactic models, consistency checks can be carried out and such samples could lead to exciting and unexpected discoveries. ",
        "introduction": "Gravitational microlensing is now established as a powerful tool for various aspects of Galactic science. The concept of gravitational microlensing, where a background object is temporarily brightened by the gravitational lensing effect of a passing foreground object, is believed to have been initially conceived by Einstein around 1912 \\citep{ren97}. As he stated in a later publication \\citep{ein36}, his opinion was that \"there is no great chance of observing this phenomenon\". Fortunately the idea was later resurrected by several authors \\citep{lie64,ref64}, in particular \\citet{pac86}. Following the proposal of Paczy\\'nski, many experiments began to hunt for microlensing events towards the Galactic bulge and the Magellanic Clouds: OGLE \\citep{uda92}, EROS \\citep{aub93}, MACHO \\citep{alc93}, MOA \\citep{bon01}. To date thousands of microlensing events have been detected toward several directions, including the Galactic bulge \\citep{tho05,sum06,ham06}, Magellanic Clouds \\citep{tis07,wyr09,wyr10}, and M31 \\citep{cal05,dej06} and even in surveys of bright nearby stars \\citep{fuk07,gau08}.  The purpose of this paper is to deal with a previously under-exploited set of targets, namely quasars. Although quasar microlensing has been observed for many years \\citep[e.g.][]{cha79,irw89,ang08} this has, with a few exceptions \\citep[e.g.][]{bya70}, almost exclusively been lensing of background quasars by compact objects in the intervening (strong) lensing galaxy. In contrast, we are interested in the lensing of background quasars by stars in the Milky Way. This is a particularly prescient subject because imminent large-scale optical surveys, such as Pan-STARRS \\citep{kai02} and LSST \\citep{tys02,ive08}, will monitor the whole available sky every few nights down to magnitudes which are sufficient to monitor several million background quasars. If large enough samples of such microlensing events could be accrued then it will be possible to use them to constrain models of the Galaxy.  One reason why quasars are potentially important targets is that they break a number of the degeneracies known to inhibit the analysis of microlensing events from sources within our own Galaxy. In general, for lensing of Galactic sources it is difficult to determine the properties of the lensing population because all of the information about the event is contained in one observable, namely the event time-scale. As a consequence, there are a number of degeneracies which need to be broken in order to determine, for example, the lens mass or distance. For sources within the Milky Way, in general, we do not know the distance or transverse velocity of the star which is lensed. For quasars the interpretation is made significantly easier because we can assume that the source is at infinity and that the transverse velocity is negligible. In this case the lens mass is simply a function of the (known) time-scale and the (unknown) lens distance and velocity, \\begin{equation} \\label{equ1} M \\approx 0.3\\rm \\sm\\left({t_{\\rm E} \\over 30~ \\rm day}\\right)^2 \\left({\\dlens \\over 4~ \\rm kpc}\\right) \\left({\\mu_t \\over 9.4~ \\rm mas~yr^{-1}}\\right)^2, \\end{equation} where $\\dlens$ and $\\mu_t$ are the lens distance and proper motion, respectively, and $t_{\\rm E}$ is the event time-scale. This makes the interpretation of quasar microlensing events significantly easier than that of Galactic sources and opens up the possibility of carrying out model-independent mass determinations using high precision astrometry from, for example, ESA's Gaia mission \\citep{per01}. Even without a measurement of the lens distance/velocity, it could also be possible to use higher-order microlensing effects such as the microlensing parallax \\citep{gou92} to break the degeneracy. Clearly quasar microlensing has the potential to become an important avenue for probing both stellar and Galactic astrophysics.  In this paper we carry out Monte Carlo simulations to address this issue, presenting an all-sky prediction of the rate and optical depth of microlensed quasars. The paper is organised as follows. In Section \\ref{sec:model}, we describe the model and theory. In Section \\ref{sec:results}, we present the main results, dealing with the time-scale distribution (\\ref{sec:tE}), the predictions for the optical depth and event rate (\\ref{sec:event_rate}) and also the prospects for determining the lens masses (\\ref{sec:lens_mass}). We present a summary and discussion in Section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In this paper we have carried out Monte Carlo simulations to make predictions for the microlensing of quasars by stars (and stellar remnants) in the Milky Way. Using a simple model of the Galaxy we have predicted the all-sky microlensing optical depth and have calculated the time-scale distributions and event rates for future large-area sky surveys, both for $i$- and $z$-band. We found that the total event rate increases rapidly with increasing magnitude limit, and predict that around 10 to 20 events will be detected per year from surveys like Pan-STARRS and LSST (see Section \\ref{sec:event_rate}). Most events are caused by lenses in the thin disc, even though extinction hampers the detectability of low-latitude events. If we could identify and monitor large numbers of quasars at low latitudes \\citep[e.g.][]{lee08}, then it may even be possible to increase the number of detected events.  We have also highlighted how such quasar microlensing can be used to carry out robust and model-independent lens mass determinations, predicting that with accurate astrometry from Gaia or SIM-lite it will be possible to accurately recover the lens mass for around one event per year. Microlensing is uniquely placed to provide mass measurements for individual stars (as opposed to those in binary systems) and as such will undoubtedly play a crucial role in stellar astronomy.  Although these results are very promising, there are a couple of technical challenges which must be overcome in order to detect these microlensing events. In this paper we have assumed a detection efficiency of 100 per cent, which we have justified by noting that practically all events should be reasonably well-sampled by surveys like LSST. However, there are practical difficulties which will potentially reduce the efficiency.  Firstly, it is well-known that quasars themselves are intrinsically variable \\citep[e.g.][]{koz10}. They can exhibit both long-term low-level variability as well as outburst phenomena. Indeed, one recent paper claimed to have detected a spectacular UV flare in a quasar behind M31 \\citep{meu10}, although the authors admit that an alternative explanation could be microlensing. Some authors have claimed that nearly all quasars are microlensed by intervening (cosmological) compact objects \\citep[e.g.][]{haw96}, although more recent studies have challenged this hypothesis \\citep[e.g.][]{dev05}. This latter paper analyses a variability diagnostic (the structure function) of around 40,000 quasars, concluding that the most-likely cause of quasar variability is flares due to disc instabilities. They remark that the variability is, in general, both chromatic and asymmetric. This is an important issue since the microlensing events which we are considering in this paper should be both achromatic and symmetric. Therefore one avenue for disentangling microlensing from intrinsic variability would be to utilise the multi-band photometry from surveys such as LSST. In addition, the amplitude of the quasar variability is generally at a lower level than our microlensing amplitude (e.g. \\citealp{ses07} show that the rms scatter for quasars is $\\sim 0.07$ mag in the $g$-band and, because the rms decreases for longer wavelengths, the variatbility will be even less for microlensing surveys carried out in $i$- or $z$-bands).  Although most microlensing events are symmetric and achromatic, this is not always the case. There are higher-order microlensing effects which can cause deviations from the standard symmetric and chromatic form, such as binarity in the lens or finite-source effects. Both of these will probably be of only limited importance; for standard galactic microlensing the fraction of events which show clear binary signatures is of the order 3 per cent \\citep[e.g.][and references therein]{sko08}, while for finite-source effects it is even less. Although it is safe to assume that the fraction of binary lens events for quasar sources will be similar to that for bulge sources, there is no reason to assume this will also hold for finite-source effects. If we assume that a typical quasar accretion disc is of the order $10^{13}$ cm, then the angular size will be less than a few $\\mu$as. This angular size is comparable to a giant star in the Galactic bulge, which means that the finite-source effect for quasar microlensing should be no-more important than for standard Galactic microlensing; for these events finite-source effects only come into play when the amplification reaches around one hundred or more \\citep[e.g.][]{cas06}. Therefore for the vast majority of quasar microlensing events the light-curves will indeed be symmetric and chromatic. For the limited number of events for which the magnification is sufficient to detect finite-source signatures, microlensing will enable constraints to be placed on the profiles of quasar accretion discs, although this could only be done in very fortuitous circumstances (e.g. when a low-redshift quasar is very highly magnified).  A second obstacle to detecting microlensing of quasars is the difficulty of identifying the quasars in the presence of a luminous lens. The quasars will be identified via cuts in colour-colour space, and so if the lens is sufficiently bright it could compromise the quasar selection. Fortunately in most cases the lenses will be significantly fainter than the sources (see Fig. \\ref{fig:lens_luminosity_marg}). We calculate that for only a quarter of LSST events will the lens be brighter than the source and for just over half of all events the lenses will be fainter than the single-epoch magnitude limit of $i=24$. However, this issue will be particularly important for the events from which we wish to determine the lens mass, since for these we $require$ the lens to be luminous so that its distance and proper motion can be determined. For these events it will be important to spectroscopically confirm the presence of a quasar behind the lens, or to obtain high-resolution imaging a number of years after the event so that the lens and quasar can be resolved.  In contrast to the above challenges that hamper the detection of quasar microlensing events, there are avenues which could enhance our ability to detect events. If we could build up an accurate astrometric map of the foreground lenses and their proper motion (using, for example, data from the Gaia mission) then it could be possible to $predict$ which quasars will be lensed in advance of the amplification. This method has been suggested as an approach to detect lensing by pulsars \\citep{dai10} but it is equally applicable to any lens with an accurate proper motion measurment. Furthermore, this could be applied to fainter quasars, such as those provided by the complete co-added LSST survey, which will reach a depth of around $i=27$. At this magnitude the number density of quasars climbs to around 3000 per square degree (see Fig. \\ref{fig2}.) Another avenue is that of astrometric microlensing where, instead of detecting the photometric lensing signatures, we identify the astrometric shift of the image centroid as the quasar is lensed \\citep[e.g.][]{hog95,wal95}. This will be a demanding task as it requires very high precision astrometry for our faint source population; the maximum astrometric shift is 0.354 times the Einstein radius \\citep{dom00} and for our events typical Einstein radii are around 1 mas. As above we could envisage targeting specific quasars based on predictions from lens proper motions, but this would still require astrometry at a level of at least tens of $\\mu$as.  In this work we have only considered lenses within our own Galaxy, yet one might wonder whether quasars could be lensed by stellar mass objects at cosmological distances. However, this is a remote possibility as it would require a very high density of such objects and even then the time-scales for such events are a couple of orders of magnitude longer than what we find for Galactic lenses \\citep[see, for example,][]{wam06}]. The notable exception to this, as noted in the Introduction, is the phenomena of microlensing of macro-lensed (i.e. multiple-imaged) quasars by stellar mass lenses in the intervening lensing galaxy. Here the density of lenses is sufficient to produce an optical depth of order unity and hence detecting microlensing becomes feasible. Surveys such as LSST and Pan-STARRS will have an immense impact on this field as they will monitor the light-curves for many thousands of multiply-imaged quasars.  Although we are unlikely to find events from lenses at cosmological distances (with the exception of macro-lensed quasars), it may be possible to detect lenses from populations just outside our own Galaxy. For example, one could identify quasars behind the LMC \\citep[e.g.][]{koz09} or M31 \\citep{huo10} and study their lensing by foreground stars in these galaxies. As we mentioned above, there may have already been an event of this kind detected in a quasar behind M31 \\citep{meu10}.  In conclusion, we have shown that imminent ground-based surveys will be able to detect a number of quasar microlensing events every year. The efforts of Pan-STARRS and LSST will, after five years of operation, detect up to one hundred events. Although this is still too few events to constrain Galactic models due to the overwhelming numbers of parameters, it will allow for valuable consistency checks to be carried out on these models and could lead to exciting and unexpected discoveries."
    },
    "1005/1005.0359_arXiv.txt": {
        "abstract": "{Modern radio interferometers sensitive to low frequencies will make use of wide-band detectors (with bandwidths of the order of the observing frequency) and correlators with high data processing rates. It will be possible to simultaneously correlate data from many sub-bands spread through the whole bandwidth of the detectors. For such wide bandwidths, dispersive effects from the atmosphere introduce variations in the fringe delay which change through the whole band of the receivers. These undesired dispersive effects must be estimated and calibrated with the highest precision.} {We studied the achievable precision in the estimate of the ionospheric dispersion and the dynamic range of the correlated fringes for different distributions of sub-bands in low-frequency and wide-band interferometric observations. Our study is focused on the case of sub-bands with a bandwidth much narrower than that of the total covered spectrum (case of LOFAR).} {We computed the formal statistical uncertainty of the ionospheric delay, the delay ambiguity and the dynamic range of the correlated fringes using four different kinds of distributions of the sub-bands: constant spacing between sub-bands, random spacings, spacings based on a power-law distribution, and spacings based on Golomb rulers (sets of integers, $n_i$, whose sets of differences, $n_j-n_i$, have non-repeated elements).} {We compare the formal uncertainties in the estimate of ionospheric effects in the data, the ambiguity of the delays, and the dynamic range of the correlated fringes for the four different kinds of sub-band distributions.} {For a large number of sub-bands ($> 20$, depending on the delay window) spacings based on Golomb rulers give the most precise estimates of dispersive effects and the highest dynamic ranges of the fringes. Spacings based on the power-law distribution give similar (but slightly worse) results, although the results are better than those from the Golomb rulers for a smaller numbers of sub-bands. Random distributions of sub-bands result in relatively large dynamic ranges of the fringes, but the estimate of dispersive effects through the band is worse. A constant spacing of sub-bands results in very bad dynamic ranges of the fringes, but the estimates of dispersive effects have a precision similar to that obtained with the power-law distribution. Combining all the results, the power-law distribution gives the best compromise between homogeneity in the sampling of the bandwidth, precision in the estimate of the ionospheric dispersive effects, dynamic range of the correlated fringes, and ambiguity of the group delay.} ",
        "introduction": "Modern radio interferometers sensitive to low frequencies, like the LOw Frequency ARray (LOFAR, see, e.g., de Vos, Gunst, \\& Nijboer \\cite{Vos2009} and references therein), will make use of wide-band detectors (with a bandwidth of the order of the observing frequency) and correlators with a high data processing rate. For the case of LOFAR, it will be possible to correlate a bandwidth of 42\\,MHz, which can be divided into many sub-bands to be spread through a wide band (either 10--90\\,MHz, 110--190\\,MHz, 170--230\\,MHz, or 210--250\\,MHz; de Vos, Gunst, \\& Nijboer \\cite{Vos2009}). Therefore, the correlator will be able to simultaneously cover a large portion of the low-frequency spectrum. However, atmospheric dispersion may introduce strong frequency-dependent effects, which will be so different for each sub-band and must be estimated and calibrated with care. At low frequencies, the ionosphere is the main limiting factor in the quality of the interferometric observations. For the calibration of the ionospheric dispersion, the Total Electron Content (TEC) must be estimated through the field of view (FoV) over each antenna and for each time. The TEC can be estimated, for instance, from Global-Positioning-System (GPS) data (e.g. Ros et al. \\cite{Ros2000}; Todorova, Hobiger, \\& Schuh \\cite{Todorova2006}). With these techniques, however, it is not possible to obtain high-resolution estimates of the ionospheric turbulent screen over the FoV. For that purpose, it is necessary to apply self-calibration strategies, using therefore the interferometric data to derive the structure of the TEC screen over each station and its evolution (see Intema et al. \\cite{Intema} and Cohen \\& Rottgering \\cite{Cohen} for a discussion on the ionospheric screens).  The distribution of sub-bands through the whole bandwidth of the detectors affects the scientific information that can be extracted from the observations, but is also important for a precise and accurate estimate of the atmospheric dispersion from the interferometric observables. The distribution of sub-bands has also an effect in delay space, since the interferometric fringes have a shape related to the Fourier transform of the bandpass (i.e., in our case, the distribution of sub-bands). Therefore, we should distribute the sub-bands in such a way that the spectral coverage and sampling are maximized, but the dynamic range of the fringes (i.e., the height of the fringe peak relative to that of the highest sidelobe) is also maximized to improve the sensitivity of the interferometer. Finding out the right distribution of sub-bands to achieve an optimum spectral sampling and fringe dynamic range is not that simple, and the answer may depend on each case, namely, the total bandwidth to be covered, the number of sub-bands, and the bandwidth of each sub-band. For instance, Petrachenko (\\cite{Petrachenko2008}) studied the performance of ``broadband delays'', which are computed from several bands (up to 5) spread from 2$-$3\\,GHz to 11$-$14\\,GHz, as a function of the way these bands are distributed through the spectrum. Petrachenko (\\cite{Petrachenko2008}) concluded that the use of more than 2 bands, covering a total bandwidth as wide as possible, improves the performance of the interferometer.  In this paper, we report on a study of the spectral coverage, the precision in the estimate of atmospheric dispersive effects, the delay ambiguity, the dynamic range of the fringes, and the precision in the estimate of the source spectral index for different kinds of sub-band distributions and spectral configurations of an interferometer at low frequencies. The remainder of this paper is structured as follows: in Sect. \\ref{II} we describe the process of analysis. In Sect. \\ref{III} we describe the different sub-band distributions studied. In Sect. \\ref{IV} we report on the results obtained and in Sect. \\ref{V} we summarize our conclusions. ",
        "conclusions": "\\label{V}  We studied the achievable precision in the estimate of the ionospheric dispersion, the ambiguity of the group delay, the dynamic range of the correlated fringes, and the precision in the estimate of the source spectral index in low-frequency and wide-band interferometric observations for four different distributions of the sub-bands through the total bandwidth of the detectors: constant spacing between sub-bands, random spacings, spacings based on a power-law distribution, and spacings based on Golomb rulers.  For a large number of sub-bands, spacings based on Golomb rulers give the most precise estimates of dispersive effects and the highest dynamic ranges of the fringes. Spacings based on the power-law distribution give similar (but slightly worse) results, although the results are better than those from the Golomb rulers for a smaller numbers of sub-bands. Random distributions of sub-bands result in relatively large dynamic ranges of the fringes, but the estimate of dispersive effects through the band is worse. A constant spacing of the sub-bands results in very bad dynamic ranges of the fringes, but the estimates of dispersive effects have a precision similar to that obtained with the power-law distribution.  From all combinations of the number of sub-bands and the total covered bandwidth, the power-law distribution (with $\\alpha$ given by Eq. \\ref{AlphaNEq}) gives the best compromise between homogeneity in the sampling of the bandwidth, precision in the estimate of the ionsopheric dispersive effects, dynamic range of the correlated fringes, and ambiguity of the group delay. Therefore, this would be the preferable sub-band distribution to use in low-frequency (wide-band) interferometric observations.  Finally, we study how the power-law distribution that optimally samples the ionospheric dispersion is affected in the cases of 1) observing frequencies close to the plasma frequency of the ionospheric electrons, 2 chromatic effects in the structure of the sources, and 3) non-negligible noise coming from the Galaxy radiation."
    },
    "1005/1005.0029_arXiv.txt": {
        "abstract": "{\\sl Ground Level Enhancement (GLE)} events represent the most energetic class of {\\sl solar energetic particle (SEP)} events, requiring acceleration processes to boost $\\gapprox 1$ GeV ions in order to produce showers of secondary particles in the Earth's atmosphere with sufficient intensity to be detected by ground-level neutron monitors, above the background of cosmics rays. Although the association of GLE events with both solar flares and coronal mass ejections (CMEs) is undisputed, the question arises about the location of the responsible acceleration site: coronal flare reconnection sites, coronal CME shocks, or interplanetary shocks? To investigate the first possibility we explore the timing of GLE events with respect to hard X-ray production in solar flares, considering the height and magnetic topology of flares, the role of extended acceleration, and particle trapping.  We find that 50\\% (6 out of 12) of recent (non-occulted) GLE events are accelerated during the impulsive flare phase, while the remaining half are accelerated significantly later. It appears that the prompt GLE component, which is observed in virtually all GLE events according to a recent study by Vashenyuk et al., is consistent with a flare origin in the lower corona, while the delayed gradual GLE component can be produced by both, either by extended acceleration and/or trapping in flare sites, or by particles accelerated in coronal and interplanetary shocks. ",
        "introduction": "A key aspect that motivated this review is the question whether {\\sl ground level enhancement (GLE)} events, which apparently require acceleration processes that produce $\\gapprox 1$ GeV particles, originate from flare regions in the solar corona or from shocks driven by coronal mass ejections propagating through the corona and interplanetary space.  GLE events represent the largest {\\sl solar energetic particle (SEP)} events that accelerate GeV ions with sufficient intensity so that secondary particles are detected by ground-level neutron monitors above the galactic cosmic-ray background (Lopate 2006; Reames 2009b). A catalog of 70 GLE events, occurring during the last six solar cycles from 1942 to 2006, has been compiled (Cliver et al.~1982; Cliver 2006), which serves as the primary database of many GLE studies. So, GLE events are very rare, occurring only about a dozen times per solar cycle, which averages to about one event per year. While GLE events with 1 GeV energies represent the largest energies produced inside our solar system, they are at the bottom of the cosmic ray spectrum, which covers an energy range of $\\approx 10^9-10^{21}$ eV, exhibiting a ``spectral knee'' between particles accelerated inside our galaxy ($\\approx 10^9-10^{16}$ eV) and in extragalactic sources ($\\approx 10^{16}-10^{21}$ eV).  While coronal mass ejections (CMEs) are widely considered as the main drivers of geoeffective phenomena, as pointed out in the so-called ``solar flare myth'' paradigm (Gosling 1993), the acceleration site of high-energy particles detected in-situ in the heliosphere can often not unambiguously be localized, and thus we have to consider both options. Several timing studies of GLE events have shown evidence for multiple (impulsive and gradual) components, so there is often not a simple dichotomy of acceleration sites. In order to address our central question whether GeV particles producing GLE events are either of solar or of heliospheric origin, we will review various aspects of flare observations of GLE events, such as the relative timing of the release of GLE-associated particles during solar flares (Section 2.1), prompt flare-associated acceleration of GLE protons (Section 2.2), the maximum energies of gamma-ray producing particles in solar spectra (Section 2.3), the height of acceleration regions (Section 2.4), the magnetic topology of solar flare acceleration regions (Section 2.5), and the role of extended acceleration phases and particle trapping in solar flares (Section 2.6). We summarize our conclusions in Section 3. Complementary aspects of the same question from the view of CME-associated shocks are treated in the article of Gang Li, and active region characteristics of GLE-associated flares are reviewed by Nariaki Nitta in this volume.  \\begin{figure} \\centerline{\\includegraphics[width=1.0\\textwidth]{f1.eps}} \\caption{Time profiles of the GOES 1-8 \\ang\\ soft X-ray fluxes (solid red curve), GOES 0.5-4 \\ang\\ fluxes (dashed red curve), and their time derivatives (thick black curves), which are a proxi for the hard X-ray flux, and solar particle release times with uncertainties (hatched areas) of 13 GLE events during 1998-2006. The time range is centered at the hard X-ray peak time (given on the left side). The {\\sl solar particle release times (SPR)} are determined from the velocity dispersion of particles detected with WIND by Reames (2009b), corrected by +8.3 min to the photon arrival time at Earth. The start times of radio type II bursts are indicated with vertical dashed bars.} \\end{figure} ",
        "conclusions": "We explored here the question whether the largest SEP and GLE events that accelerate ions with energies of $\\gapprox 1$ GeV could be accelerated in solar flare regions, in contrast to the generally accepted paradigm of acceleration in heliospheric CME shocks. We reviewed the {\\sl pro and con} aspects from the solar flare site that are relevant to answer this question, while the complementary aspects from CME-associated shocks are discussed in the companion article by Gang Li. The major arguments for a flare-driven acceleration of GeV particles are: (1a) the coincidence of the particle release time with the impulsive flare phase (in 50\\%), (2a) the similarity of the impulsive time profiles of the first GLE component with the impulsive flare phase in hard X-rays, (3a) observational evidence of maximum energies $\\lapprox 100$ MeV in electron bremsstrahlung, and (4a) the relatively high magnetic field and free magnetic energy that is available in flare sites, compared to the locations of interplanetary CME shocks. In contrast, major arguments against this interpretation are: (1b) significant delays of the particle release time with respect to the impulsive flare phase (in 50\\%), and (2b) the lack of direct observational evidence and diagnostics for $\\gapprox 1$ GeV protonts and ions in flare sites. The conclusions are based on observations of 70 GLE events over the last six decades, in particular on the 13 GLE events during the last solar cycle 23 (1998-2006) that provided excellent new imaging data in gamma rays and hard X-rays (RHESSI), in soft X-rays and EUV (TRACE, SOHO/EIT), and particle data from IMP, WIND, and ACE. Our conclusions are:  \\begin{enumerate} \\item{Solar particle release times $t_{SPR}$ of GLE-producing particles overlap with the impulsive phase of gamma-ray and hard X-ray emission in solar flares in 6 out of 12 (non-occulted)cases, and thus the acceleration time of GLE particles is consistent with the flare site in 50\\% of the cases, taking the full duration of impulsive flare hard X-ray emission ($t_X \\approx 3-13$ min) into account.} \\item{In the remaining cases, 6 our of 12 occur delayed to the flare peak by $\\approx 10-30$ min, but observational signatures of extended acceleration and/or particle trapping are evident in all strongly delayed cases, and thus all GLE events could potentially be accelerated in flare sites. The alternative explanation of delayed second-step acceleration in CME-associated shocks cannot be ruled out, however, possibly constituting a secondary gradual GLE component.} \\item{The height of the acceleration region of \\lapprox 1 GeV electrons and ions depends on the interpretation, being $h \\lapprox 0.05$ solar radii for flare site acceleration (according to electron time-of-flight measurements), or $h \\approx 2-5$ solar radii for CME shock acceleration (according to type II source heights and GLE solar particle release delays).} \\item{The magnetic topology at the particle acceleration site is not well-known from magnetic modeling or tracing of coronal structures, but there is statistical evidence for open-field regions in most active regions that provide escape routes for GLE particles.} \\item{The recently discovered strong correlation between the spectral soft-hard-harder (SHH) evolution of solar hard X-rays and SEP events poses a new challenge. It is presently unclear how the SHH evolution can be explained in the context of the standard scenario in terms of SEP acceleration in CME-associated shocks.} \\item{The maximum particle energies observed in solar flares reach up to several 100 MeV for electrons and above 1 GeV for ions. The required particle energies are $\\lapprox 1$ GeV for the observed bremsstrahlung continuum at photon energies of $\\lapprox 10$ MeV, and $0.2-5$ GeV for pion decay radiation observed at photon energies of 10 MeV $-$3 GeV, which is sufficient to explain GLE detections.} \\item{Energies up to $\\lapprox 1$ GeV can be achieved for DC electric field acceleration in super-Dreicer fields, for stochastic acceleration, and for shock acceleration. Only DC electric field acceleration in sub-Dreicer fields can be ruled out, based on the large distances of DC fields required and the inconsistency with measured electron time-of-flight delays.} \\end{enumerate}  What do we need to answer the question of the origin of GeV particles with more certainty? The velocity dispersion measurements with IMP, WIND, and ACE helped enormously to narrow down the start window of solar particle release times, which could be complemented with directivity measurements of the particle detector IMPACT on STEREO. The most sensitive hard X-ray and gamma-ray detectors existed on CGRO, but in near future we have only RHESSI available. However, RHESSI provides excellent imaging capabilities, which revealed intriguing information on displaced electron and ion acceleration sites in the solar corona (Hurford et al.~2003, 2006). In addition we expect also improved magnetic modeling of solar flare regions that can reveal open-field and closed-field geometries with more certainty, using nonlinear force-free field (NLFFF) codes (e.g., Bobra et al.~2008; DeRosa et al.~2009) with 3-D vector magnetograph data from the {\\sl Solar Dynamics Observatory (SDO)}, possibly in conjuction with stereoscopic 3-D reconstruction of flaring active regions using {\\sl STEREO/EUVI}.  In conclusion, acceleration of GeV particles in flare sites is a possibility that cannot be firmly ruled out with the current localization capabilities of energetic particles. Certainly we have evidence for both acceleration in coronal flare sites and in heliospheric CME shocks, often appearing concomitantly, but with different (impulsive vs. gradual) time scales, relative timing, and charge state characteristics. While one-sided emphasis has been given to both, either flares (the ``big flare syndrome''), or CMEs (the ``flare myth''; Gosling 1993), there is a consensus now that both flare and CME phenomena are part of a common magnetic instability, and that both are being able to accelerate particles to high energies. The remaining questions are then mostly what the relative proportions of both components are and how we can discriminate between them. A preliminary answer is that the observations are mostly consistent with a flare-associated ``prompt GLE component'' and a CME-associated ``delayed GLE component''."
    },
    "1005/1005.1219_arXiv.txt": {
        "abstract": "Non-linear CMB temperature anisotropies up to the third-order on large scales are calculated. On large scales and in the Sachs-Wolfe limit, we give the explicit expression for the observed temperature anisotropy in terms of the primordial curvature perturbation up to third-order. We derived the final bispectrum and trispectrum of anisotropies and the corresponding non-linear parameters, in which the contributions to the observed non-Gaussianity from primordial perturbations and from the non-linear mapping from primordial curvature perturbation to the temperature anisotropy are transparently separated. ",
        "introduction": "In the past few years, extensive attention has been attracted to the investigation of cosmological perturbations \\cite{Lyth:1998xn} beyond the linear order. The importance of studying  non-linear perturbations comes two aspects. Firstly, forthcoming experiments of Cosmic Microwave Background (CMB) and Large-scale Structure (LSS) will be able to detect the non-linear structures in these perturbations. The observational detection of the non-linearities through the statistical non-Gaussianity (see \\cite{Bartolo:2004if} for a comprehensive review and \\cite{Komatsu:2010hc} for a review of recent observational progress ) of perturbations has become one the the primary targets of the cosmology. On the other hand, non-linearities, which encode the \\emph{interactions} in the early universe, would definitely bring us new understandings of both the early universe and the fundamental physics.  A large amount of efforts have been devoted to the calculation of the statistics of curvature perturbation $\\zeta$, like primordial bispectrum and trispectrum on large scales, pioneered by Maldacena \\cite{Maldacena:2002vr}{\\footnote{See e.g. \\cite{Koyama:2010xj} for a recent short review of methods and techniques in calculating and analyzing primordial non-Gaussianities from inflationary models and references therein.}}. However, these are not the  observed non-Gaussianities of (e.g.) CMB temperature anisotropies $\\frac{\\Delta T}{T}$. Conventionally, one may use the linear order relation $\\frac{\\Delta T}{T} = -\\frac{1}{3}\\Phi = -\\frac{1}{5}\\zeta$ to evaluate the angular bispectrum or trispectrum of CMB, assuming the contributions from the second-order or secondary effects are negligible comparing to the primordial ones \\cite{Komatsu:2010fb,Komatsu:2010hc,Komatsu:2001rj}. However, in light of increasingly precise observations, a full treatment of the higher-order radiation transfer functions of the CMB anisotropies is needed, which will allow us to make definite prediction of CMB non-Gaussianities (see \\cite{Bartolo:2010qu} for a recent review for non-Gaussianity on the CMB).  The research on non-linear temperature anisotropy due to gravitational perturbations was pioneered by \\cite{Pyne:1995bs,Mollerach:1997up,Matarrese:1997ay}, in which the second-order generalization of Sachs-Wolfe (SW) effect and Integrated Sachs-Wolfe (ISW) effect were derived. These results were extended in \\cite{Bartolo:2005kv,Bartolo:2004ty,Tomita:2005et,Tomita:2007db,Boubekeur:2009uk} where the second-order radiation transfer function on large scales was calculated, and in \\cite{D'Amico:2007iw} where the general expression for anisotropy due to gravitational perturbations up to the third-order and in \\cite{Bartolo:2005fp} where an elegant and non-perturbative non-linear anisotropy were got. The non-linear anisotropies have also been analyzed in \\cite{Maartens:1998xg,Dunsby:1997xy}  base on the covariant approach to cosmological perturbations (see \\cite{Langlois:2010vx} for a recent review). Various secondary contributions to the non-Gaussianities have also been extensively studied, including the weak gravitational lensing and its correlation with ISW effect \\cite{Seljak:1998nu,Goldberg:1999xm,Serra:2008wc,Hanson:2009kg,Mangilli:2009dr,Creminelli:2004pv,Cooray:2002ee}, which is expected to be the dominant contamination of $f_{\\textrm{NL}}^{\\textrm{local}}$, correlation between lensing and Sunyaev-Zel'dovich (SZ) effect \\cite{Goldberg:1999xm}, inhomogeneous recombination \\cite{Khatri:2008kb,Senatore:2008vi,Senatore:2008wk}, small-scale dark matter clustering \\cite{Bartolo:2008sg,Pitrou:2008ak}. A systematic treatment of the transfer function on all scales, which involves solving  the full Boltzmann equations through the recombination phase and then from the surface of last-scattering until today, has been performed in \\cite{Bartolo:2006cu,Bartolo:2006fj,Nitta:2009jp,Pitrou:2008ak,Pitrou:2008hy,Pitrou:2008ut,Pitrou:2010sn,Beneke:2010eg} at second-order.  In this note, we calculate the CMB temperature anisotropy up to the third-order in primordial curvature perturbation $\\zeta$, which can be viewed as non-linear generalization of linear-order relation $\\Delta T/T = -\\zeta/5$. We follow the same strategy in \\cite{Pyne:1995bs,Mollerach:1997up,Matarrese:1997ay,D'Amico:2007iw,Boubekeur:2009uk}. First, we calculate the gravitational redshift of a given photon from the emission surface to today's observer, which will give the observed anisotropy in terms of metric perturbations $\\frac{\\Delta T}{T} = \\frac{\\Delta T}{T}[\\Phi,\\Psi,\\cdots]$. Then by using the conservation of curvature perturbation $\\zeta$, in the large-scale limit, we determine the initial conditions for the metric perturbations in matter-dominated era, more precisely,  the values $\\Phi_e = \\Phi_e[\\zeta]$ and $\\Psi_e= \\Psi_e[\\zeta]$ on the emission surface. Combining these two procedures will give the final non-linear mapping from $\\zeta$ to $\\frac{\\Delta T}{T}$: $\\frac{\\Delta T}{T}=\\frac{\\Delta T}{T}[\\zeta]$. Together with previous results of primordial non-Gaussianities of $\\zeta$  got in the literature, our formalism is ready to make prediction of the final observed CMB non-Gaussianity. Obviously, since we match the metric perturbations in matter era directly to those on the emission surface after last-scattering, our formalism includes only the gravitational redshift of photons, without considering the dynamics of photon-baryon plasma. Thus our result is valid only for large-scale anisotropies, which enter the horizon after decoupling.      This note is organized as follows. In the following section, we describe the temperature anisotropy induced by the gravitational perturbations from the emission surface to the observers. We give the general expression for the temperature anisotropy up to the third-order in metric perturbations. In the next section, by using the conserved curvature perturbation and taking the large-scale limit, we determine the initial conditions for metric perturbations in matter-dominated era. Then we give the non-linear mapping from primordial curvature perturbation to the temperature anisotropy, in the Sachs-Wolfe limit. Finally we give a short conclusion. ",
        "conclusions": "In this note, following the approach developed in \\cite{Pyne:1995bs,Mollerach:1997up,Matarrese:1997ay}, the general expression for the observed CMB anisotropy up to the third-order is calculated \\cite{D'Amico:2007iw}. In the Sachs-Wolfe limit, we derive the non-linear relation between the observed  anisotropy and the conserved primordial curvature perturbation $\\zeta$, up to the third-order in $\\zeta$. (\\ref{DeltaT2})-(\\ref{DeltaT3}) can be viewed as non-linear generalization of familiar linear relation $\\Delta T/T= -\\zeta/5$. Our formalism is valid for large-scale anisotropies, which re-enter the horizon after decoupling. The results (\\ref{bispectrum})-(\\ref{trispectrum}) clearly show the different contributions to the observed non-Gaussianity from primordial non-Gaussianities in $\\zeta$ and non-linear mapping from $\\zeta$ to $\\frac{\\Delta T}{T}$ due to gravitational perturbations. We also derive the non-linear parameters $f_{\\textrm{NL}}$ and $g_{\\textrm{NL}}$ (eq.(\\ref{f_NL_final})-(\\ref{g_NL_final})), which enter the theoretical predictions for the angular bispectrum and trispectrum of CMB respectively.  We do not expect the non-linear mapping in the SW limit (\\ref{DeltaT2})-(\\ref{DeltaT3}) would give a major contribution to the observed non-Gaussianities, especially comparing scenarios where large ``primordial\" non-Gaussianities can be generated (see \\cite{Koyama:2010xj} for a recent review and references therein). However, the other secondary anisotropies, especially the correlation between lensing and ISW effect which we do not discuss in this note, is expected to give contribution to the final non-Gaussianity \\cite{Seljak:1998nu,Goldberg:1999xm,Serra:2008wc,Hanson:2009kg,Mangilli:2009dr,Creminelli:2004pv,Cooray:2002ee}. Another interesting issue is that, if an enhancement of non-linearity between $\\frac{\\Delta T}{T}$ and $\\zeta$ is possible, like the enhancement of primordial non-Gaussianity by small $c_s$ etc. in some inflationary scenarios. We wish to come back to these subjects in future investigations."
    },
    "1005/1005.3500_arXiv.txt": {
        "abstract": "{Dwarf irregular galaxies are relatively simple unevolved objects where it is easy to test models of galactic chemical evolution.} {We attempt to determine the star formation and gas accretion history of IC10, a local dwarf irregular for which abundance, gas, and mass determinations are available.} {We apply detailed chemical evolution models to predict the evolution of several chemical elements (He, O, N, S) and compared our predictions with the observational data. We consider additional constraints such as the present-time gas fraction, the star formation rate (SFR), and the total estimated mass of IC10. We assume a dark matter halo for this galaxy and study the development of a galactic wind. We consider different star formation regimes: bursting and continuous. We explore different wind situations: {\\em i}) normal wind, where all the gas is lost at the same rate and {\\em ii}) metal-enhanced wind, where metals produced by supernovae are preferentially lost. We study a case without wind. We vary the star formation efficiency (SFE), the wind efficiency, and the time scale of the gas infall, which are the most important parameters in our models.}{We find that only models with metal-enhanced galactic winds can reproduce the properties of IC10. The star formation must have proceeded in bursts rather than continuously and the bursts must have been less numerous than $\\sim$10 over the whole galactic lifetime. Finally, IC10 must have formed by a slow process of gas accretion with a timescale of the order of 8 Gyr.}{} ",
        "introduction": "Understanding how dwarf galaxies formed is a key aspect of modern observational cosmology because they are currently the most abundant galaxy population, and were probably even more abundant during the first epochs of the Universe. In addition, according to the hierarchical formation scenario (e.g., White \\& Frenk \\cite{white91}), they are building blocks of larger structures, hence their study is crucial to understanding galaxy formation and evolution processes. Several questions about their formation and evolution still remain unsolved, among them: {\\em i)} the origin of the differences between the irregular and spheroidal morphological types (e.g., Mateo \\cite{mateo98}), {\\em ii)} the relation between their mass and metallicity, {\\em iii)} the contribution of galactic and supernovae (SNe) winds to the evolution of the interstellar medium (ISM) metallicity, and {\\em iv)} their age and star formation history (SFH) (cf., e.g., Gavil\\'an et al. \\cite{gavilan09}). Local Universe dwarf galaxies are unique environments for studying these aspects. Owing to the intrinsic faintness of dwarf galaxies, observational constraints to their evolution, such as the chemical abundances of stellar populations of different ages, can be obtained only for nearby galaxies.  Therefore, the large amount of information available for Local Group (LG) dwarf galaxies, such as star formation rate (SFR),  \\hi\\ and H$_2$ contents, star surface density, and metallicity distribution, allows us to use them to test chemical evolution models of different kind of galaxies (cf. Carigi et al. \\cite{carigi02}; Lanfranchi \\& Matteucci \\cite{lanfranchi03}; Moll\\'a \\& D\\'\\i az \\cite{molla05}). Elemental abundance ratios and their variation over cosmic time because of ongoing star formation (SF) are among the most important constraints of chemical evolution models (e.g., Moll\\'a et al. \\cite{molla96}). In this framework, optical spectroscopy of emission-line gas in galaxies is essential for deriving chemical abundances. In particular,  \\hii~regions and planetary nebulae (PNe) provide information about the gas chemical composition at two different epochs in the galactic history: \\hii~regions provide the present time chemical abundances, while PNe probe the distant past chemistry.  Dwarf irregular galaxies (dIrrs) are of great interest because they enable detailed studies of issues such as the occurrence of galactic winds and the chemical enrichment of the interstellar and intergalactic media by the direct measurement of the metal content of their ISM. In addition, their low level of evolution, as implied by their low metallicity and high gas content, makes these systems the most similar to primeval galaxies and, therefore, the most useful for inferring the primordial galaxy conditions. It has also been proposed that dwarf irregulars could represent the local counterparts of faint blue galaxies found in excess in deep galaxy counts (e.g., Ellis \\cite{ellis97}).  In this work, we deal with the detailed chemical evolution of the dwarf irregular galaxy IC10, building several chemical evolution models constrained by the observed properties of this galaxy. Our choice is motivated by recent observations (e.g., Sanna et al.\\cite{sanna09}; Magrini \\& Gon\\c calves \\cite{magrini09}, hereafter MG09) that place strong constraints on both its SFH and its chemical content and enrichment.  We describe the baryonic content, the SFR, and the chemical abundances of IC10 in Sect.~2. The framework of the modelling is given in Sect.~3, whereas the results are presented in Sect.~4. Finally, Sect.~5 is devoted to the conclusions. ",
        "conclusions": "We have studied the chemical evolution of the dwarf irregular galaxy IC10 for which detailed abundance data are now available (MG09). We have adopted a chemical evolution model for dwarf irregulars developed by Yin et al. (\\cite{yin10}). This model includes detailed metallicity-dependent stellar yields for both massive and low and intermediate mass stars. The model also includes feedback from SNe and stellar winds and follows the development of galactic winds. We have explored three cases: {\\em i}) no winds (no feedback); {\\em ii}) normal winds, namely all the gas is lost at the same rate; and {\\em iii}) metal-enhanced winds where the metals are lost preferentially relative to H and He. We also explored several regimes of star formation: a) bursting mode with short and long bursts (a regime we called ``gasping SF''); and b) a continuous but low SFR.  We computed the evolutions of the abundances of several elements (He, N, O, S). The parameters of our model were: 1) the number of bursts; 2) the duration of bursts; 3) the efficiency of star formation; and 4) the efficiency of the galactic wind rate, assumed to be proportional to the amount of gas present at the time of the wind, which can vary from element to element in the case of the metal-enhanced wind.  The observational constraints were represented by the abundances and abundance ratios of the above-mentioned elements, by the present-time gas mass, the SFR, and the estimated total mass of IC10.  In spite of the large number of parameters, the number of observational constraints reproduced by our model makes us confident in concluding the following:   \\begin{itemize} \\item The galaxy IC10 probably formed by means of a slow gas accretion process with a long infall timescale of the order of 8 Gyr.  \\item We do not know a priori the SFH of IC10 but our numerical simulations suggest that it should have experienced between a minimum of 4 and a maximum of 10 bursts of star formation. If we were to trust the present-time SFR measured from FIR and 1.4 GHz, the duration of the bursts should be longer than the interburst period, corresponding to a gasping mode of star formation, and the SFE should not be higher than 0.4 Gyr$^{-1}$. On the other hand, if we were to assume the present-time SFR inferred from H$\\alpha$, which is higher than the other measurements, then bursts of higher SFE ($\\sim 1$ Gyr$^{-1}$) but longer interburst periods can also be acceptable. In any case, we exclude a regime of continuous SF.  \\item We exclude evolution without galactic winds. In this case, our model can reproduce neither the present-time gas fraction nor the metallicity of IC10 at the same time without including winds. Therefore, galactic winds must occur.   \\item A metal-enhanced wind is more likely than a normal one, otherwise an unreasonable amount of mass is lost by the galaxy, making it impossible to fit the present-time amount of gas. In our optimal models (e.g. models mw8, mw9, and mw12), a continuous wind develops following the first burst that carries away mostly metals, as assumed in hydrodynamical simulations. Both gasping and bursting models with metal-enhanced winds can reproduce the macroscopic properties of IC10 and the scatter in the abundance ratios.  \\end{itemize}"
    },
    "1005/1005.4927_arXiv.txt": {
        "abstract": "As deeper observations discover increasingly distant galaxies, characterizing the properties of high-redshift galaxy populations will become increasingly challenging and paramount. We present a method for measuring the clustering bias of high-redshift galaxies from the field-to-field scatter in their number densities induced by cosmic variance.  Multiple widely-separated fields are observed to provide a large number of statistically-independent samples of the high-redshift galaxy population. The expected Poisson uncertainty is removed from the measured dispersion in the distribution of galaxy number counts observed across these many fields, leaving, on average, only the contribution to the scatter expected from cosmic variance.  With knowledge of the $\\Lambda$-Cold Dark Matter power spectrum, the galaxy bias is then calculated from the measured cosmic variance.  The results of cosmological N-body simulations can then be used to estimate the halo mass associated with the measured bias. We use Monte Carlo simulations to demonstrate that Hubble Space Telescope pure parallel programs will be able to determine galaxy bias at $z\\gtrsim6$ using this method, complementing future measurements from correlation functions. ",
        "introduction": "Galaxy bias, the spatial clustering strength of galaxies relative to the matter field, encodes important information about the connection between the luminous baryons of galaxies and the dark matter (DM) halos that host them.  The galaxy bias $b$, defined as the ratio between the galaxy and matter correlation functions (CFs) $b^{2}\\equiv \\xi_{gg}/\\xi_{mm}$, reflects both the clustering of galaxy-scale DM halos and how luminous galaxies populate halos of a given mass. A self-consistent connection between the luminosity function (LF), angular CF, and mass-to-light ratio of galaxies can be developed \\citep[e.g.,][]{van_den_bosch2003a,vale2004a,conroy2006a,lee2009a}.  CF estimates of bias, which require large samples of galaxies spanning a wide range of angular separations, are remarkably successful at low-redshift \\citep[e.g.,][]{norberg2001a,zehavi2002a}. Angular CF measurements on smaller samples have been performed for Lyman Break Galaxies (LBGs) \\citep[][]{giavalisco1998a,arnouts1999a,ouchi2004a,ouchi2005a,foucaud2003a,hamana2004a,adelberger2005a,allen2005a,kashikawa2006a,lee2006a,lee2009a} and Lyman-$\\alpha$ emitters \\citep{ouchi2003a,hamana2004a,kovac2007a,gawiser2007a} over a range of redshifts. These studies have extended high-significance $(>3\\sigma)$ detections of bias out to $z\\sim5$.    At the highest redshifts ($z\\gtrsim5.5$), limited samples have made significant detections of galaxy bias difficult.  \\citet[][]{mclure2009a} used a $0.63$deg$^{2}$ area observed by the UKIRT Infrared Deep Sky Survey Ultra Deep Survey and the Subaru XMM-Newton Survey to measure the clustering of $5<z<6$ LBGs with $z<26$AB ($b=5.2_{-0.8}^{+1.2}$), but most of their systems lie at redshifts $z<5.5$. \\citet[][]{overzier2006a} reported a measurement of bias for $i$-dropout galaxies in Great Observatories Origins Deep Survey (GOODS) ($b=4.1_{-2.6}^{+1.5}$) but no detection of clustering in the Ultra Deep Field. A study of $z$-dropout galaxies in the Subaru Deep Field and GOODS-N by \\citet{ouchi2009a} indicated that $z\\sim7$ galaxies may be strongly clustered, but measuring bias from only 22 candidates is challenging.  Given the difficulty of measuring clustering of high-redshift galaxies from the angular CF, an exploration of complementary approaches is warranted. As a salient example, \\citet{adelberger1998a} used a counts-in-cells analysis \\citep[][]{peebles1980a} to measure the clustering of redshift $z\\sim3$ LBGs from statistical fluctuations in their number densities. Having obtained spectroscopic redshifts for $N\\sim270$ LBGs in six fields, they split their samples into redshift slices and calculated the variance in galaxy counts in rectangular volumes. The bias was then measured from the excess dispersion in the number counts beyond the Poisson variance.  In this {\\it Letter}, we present calculations that demonstrate the bias of high-redshift ($z\\gtrsim6$) galaxies can also be measured from field-to-field dispersion in galaxy counts from photometric dropout samples. Specifically, we use Monte Carlo simulations to demonstrate that pure parallel programs with the Hubble Space Telescope (HST) that obtain large numbers ($\\gtrsim100$) of statistically-independent samples of $z\\approx6-7$ galaxies with the Advanced Camera for Surveys (ACS) or Wide Field Camera 3 (WFC3) \\citep[e.g.,][]{ratnatunga2002a,sparks2002a,trenti2008b,yan2008a} can provide a significant measurement of high-redshift galaxy bias. Throughout, we adopt the cosmological parameter values determined by the $5$-year \\WMAP~joint analysis \\citep{komatsu2009a}. ",
        "conclusions": "We have presented a simple counts-in-cells method \\citep[e.g.,][]{peebles1980a,adelberger1998a} for measuring the bias of high-redshift galaxies from the field-to-field variation in number counts induced by cosmic variance.  The number of high-redshift galaxies of a given luminosity are measured in a large number independent, widely-separated fields.  The Poisson contribution to the variance in the number counts across these fields is removed, leaving a cosmic variance contribution that depends on the bias and the mean matter overdensity fluctuations on the scale of the survey field volume. We use Monte Carlo simulations of the bias measurement to estimate the effectiveness of the method for example galaxy populations, $i$-dropouts at $z\\sim6$ in ACS-sized fields and $\\FSIXLP$-dropouts at $z\\sim7$ in WFC3-sized fields. Our method can be adapted for other dropout filter strategies. At $z\\sim6$ ($z\\sim7$) the method should provide a measurement of galaxy bias with $SNR\\gtrsim3$ for luminosities near $\\MUV\\approx\\Mstar$ when $\\Nfields\\ge40$ ($\\Nfields\\gtrsim130$) fields are used. The uncertainty on the galaxy bias improves with an increasing number of independent fields as $\\approx1/\\sqrt{\\Nfields}$. These requirements on the number of fields are concomitant with the expected number of pointings available from extensions to ongoing pure parallel HST programs \\citep[][]{trenti2008b,yan2008a}, and the potential for using this method to measure high-redshift galaxy bias is therefore promising. We also show that using correlated, nearby fields to perform the presented measurement typically leads to an underestimate of the bias owing to a correlation-induced decrease in the dispersion of galaxy number counts.  The measured bias can be associated with a characteristic mass of DM halos with similar clustering.  We show that the typical uncertainty in the bias measured using the presented method corresponds to an uncertainty in the inferred DM halo mass of $\\sim$ a few percent in $\\log M$ for $i$-dropouts.  Combining these estimates of halo mass with the observed luminosities allow for an estimate of a mass-to-light ratio, which can constrain the connection between the galaxies' observed LF and their clustering \\citep[][]{yang2003a,conroy2006a}. Numerous effects can alter the correspondence between bias and halo mass, and we therefore view the presented method as primarily a measure of bias and interpret the halo mass estimates with caution. Our simulations assume that the high-redshift galaxy population is essentially volume limited and pure. Methods to handle incompleteness for LBG samples have already been engineered \\citep[e.g.,][]{bouwens2008a,reddy2009a}. Contamination at a fractional level of $f_{c}$ will lower the bias by an amount $\\Deltab\\approx(1-f_{c})$ \\citep[e.g.,][]{ouchi2004a}. Using simulations that include red $z\\sim2$ galaxy interlopers with $f_{c}=0.2$ and a bias of $b\\approx5.5$ \\citep{quadri2007a}, we find that the measured bias at $z\\sim6$ is lowered by $\\sim17\\%$. This systematic error is comparable to the statistical error on the bias when $\\Nfields\\gtrsim200$. Lastly, not every parallel field will be useful for measuring high-redshift galaxy counts owing, e.g., to possible bright star or Galactic reddening contamination, and the efficiency of the presented method may be correspondingly decreased."
    },
    "1005/1005.1505_arXiv.txt": {
        "abstract": "{103P/Hartley~2 was selected as the target comet for the \\textit{Deep Impact} extended mission, \\textit{EPOXI}, in October 2007. There have been no direct optical observations of the nucleus of this comet, as it has always been highly active when previously observed.} {We aimed to recover the comet near to aphelion, to a) confirm that it had not broken up and was in the predicted position, b) to provide astrometry and brightness information for mission planning, and c) to continue the characterisation of the nucleus.} {We observed the comet at heliocentric distances between 5.7 and 5.5 AU, using FORS2 at the VLT, at 4 epochs between May and July 2008. We performed $VRI$ photometry on deep stacked images to look for activity and measure the absolute magnitude and therefore estimate the size of the nucleus.} {We recovered the comet near the expected position, with a magnitude of $m_R = 23.74\\pm0.06$ at the first epoch. The comet had no visible coma, although comparison of the profile with a stellar one showed that there was faint activity, or possibly a contribution to the flux from the dust trail from previous activity. This activity appears to fade at further epochs, implying that this is a continuation of activity past aphelion from the previous apparition rather than an early start to activity before the next perihelion. Our data imply a nucleus radius of $\\le 1$ km for an assumed 4\\% albedo; we \\emph{estimate} a $\\sim 6\\%$ albedo. We measure a colour of $(V-R)=0. 26\\pm0.09$.} {} ",
        "introduction": "The \\textit{EPOXI} mission\\footnote{\\texttt{http://epoxi.umd.edu}} is a combination of two missions that reuse the surviving fly-by part of the \\textit{Deep Impact} (DI) spacecraft. These missions were previously known as \\textit{DIXI} (Deep Impact eXtended Investigation) and \\textit{EPOCh} (Extrasolar Planet Observation and Characterisation), both of which were designed to use the DI high resolution camera, but with quite different scientific goals. The first was a direct extension of the original DI mission to Comet 9P/Tempel~1 \\citep{Ahearn05}, aimed at furthering our understanding of the nuclei of Jupiter Family Comets (JFCs) by flying on to a second comet and imaging its nucleus in detail. The second part of the mission (EPOCh) makes use of the high resolution camera as a small aperture space telescope to study transiting extrasolar planets while the spacecraft is enroute to its comet target.  Following the DI encounter with 9P/Tempel~1 the spacecraft was left on an orbit that gave the possibility of redirecting it to one of two JFCs; 85P/Boethin or 103P/Hartley~2. Comet Boethin was the preferred target as it would mean a shorter mission, however it has not been observed since 1986 and could not be recovered despite a deep search over a large area of sky covering the entire orbital position error ellipse \\citep{Meech10}. Therefore Comet Hartley~2 was selected as the target in October 2007, and the DI spacecraft was placed onto the trajectory that will take it to a flyby on November 4th 2010. 103P/Hartley~2 (hereafter, 103P) has the advantage that it has been regularly observed since it was discovered \\citep{Hartley1986} and has a well determined orbit. The disadvantage of this target is that, at the time it was selected, its nucleus had never been directly observed.  Previous observations of 103P have always found the comet to be highly active, even at relatively large heliocentric distance. Ground based telescopic observations were consistent with a radius $r_{\\rm N} \\le 4$ km, but place only weak limits due to the high levels of activity observed \\citep{Licandro00b,Lowry+Fitzsimmons01,Lowry03,Snodgrass06,Snodgrass08,Epifani08}. The previous observations are summarised in table \\ref{prev_obs}. The previous best measurement of the nucleus radius comes from \\citet{Groussin04}, who observed the comet when it was highly active at $R_{\\rm h} = 1.2$ AU, using the \\textit{Infrared Space Observatory} (ISO). They applied a coma subtraction technique previously used by this group with HST data \\citep[e.g.][]{Lamy99}, where the high spatial resolution possible using space based (diffraction limited) observations of nearby comets allows separation of the nucleus signal from the surrounding coma. Using this technique, these authors found a nucleus radius of $0.71\\pm0.13$ km for 103P. Interestingly, for this nucleus to produce the high level of activity seen in 103P, 100\\% of the surface must be active, which goes against the general tendency of JFCs to have small fractional active areas \\citep{Ahearn95}. \\citet{Lisse09} obtained images of the comet using the \\textit{Spitzer} space telescope. They also find a small nucleus with a large active area; their results are discussed further in section \\ref{discussion}.   \\begin{table} \\begin{minipage}[t]{\\columnwidth} \\caption{Previous observations and nucleus size measurement for 103P.} \\label{prev_obs} \\centering \\renewcommand{\\footnoterule}{}  % \\begin{tabular}{ccccccl} \\hline \\hline Date & $R_{\\rm h}$\\footnote{Heliocentric distance (AU): All post-perihelion except for L09.} & $\\Delta$\\footnote{Geocentric distance (AU).} & $\\alpha$\\footnote{Solar phase angle (degrees).} & $m_R$\\footnote{$R$-band magnitude.} & $r_{\\rm N}$\\footnote{Nucleus radius (or upper limit) in km assuming 4\\% albedo and phase function of 0.035 mag deg$^{-1}$} & ref.\\footnote{L00: \\citet{Licandro00b}; LF01: \\citet{Lowry+Fitzsimmons01}; L03: \\citet{Lowry03}; S06, S08: \\citet{Snodgrass06,Snodgrass08}; ME08: \\citet{Epifani08}; G04: \\citet{Groussin04}; L09: \\citet{Lisse09}} \\\\ \\hline \\multicolumn{6}{l}{Ground-based optical:}\\\\ 28/03/93 & 4.7 & 3.8 & 3.3 & 20.1 & $\\le 5.3$ & L00 \\\\ 10/12/98 & 3.6 & 3.9 & 14.3 & 19.0 & $\\le 6.4$ & LF01 \\\\ 08/06/99 & 4.6 & 4.2 & 12.5 & 20.1 & $\\le 6.3$ & L03 \\\\ 06/03/05 & 3.2 & 2.3 & 5.0 & 18.7 & $\\le 4.1$ & S06 \\\\ 03/04/05 & 3.4 & 2.5 & 7.9 & 19.1 & $\\le 4.3$ & ME08 \\\\ 01/03/06 & 5.0 & 4.3 & 8.3 & 21.0 & $\\le 4.5$ & S08 \\\\ \\hline \\multicolumn{6}{l}{Space-based infrared:}\\\\ 05/02/98 & 1.2 & 0.9 & 53.2 & -- & $0.71\\pm0.13$ & G04\\footnote{Based on a coma subtraction technique.} \\\\ 12/08/08 & 5.5 & 4.9 & 9.5 & -- & $0.57\\pm0.08$ & L09\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table} ",
        "conclusions": "We observed comet 103P/Hartley~2 using FORS2 at the VLT on 4 epochs while the comet was beyond 5.5 AU from the Sun, post-aphelion. We find:  \\begin{enumerate} \\item The comet was not visibly active in combined images at each epoch, however comparison of the profile with the image PSF revealed that there was some near nucleus dust present, even at this large distance pre-perihelion. The level of activity appeared to decrease between May and June as the comet moved inbound, implying that this was residual activity from the last active period, not the start of activity before the next perihelion passage, although we cannot rule out that this represents the trail of dust from previous perihelion passages. \\item The measured photometry puts stronger constraints on the nucleus magnitude than any previous optical observations as any activity was very weak. From the June data we find that the radius of the comet is $\\le 1$ km (assuming a 4\\% albedo). Using $r _{\\rm N}=0.57$ \\citep{Lisse09} we find that the albedo is $\\le 12\\%$. \\item The brightness decrease, and the observation of an inactive nucleus by \\citet{Lisse09} in August 2008, implies the activity ended by late July or early August, at $R_h = 5.5$ AU post-aphelion. Unfortunately our photometry from late July remains uncertain due to the comet flux being merged with background sources, but an extrapolation from May and June gives $m_R(1,1,0) \\approx 18$ for the inactive nucleus and an albedo of $\\sim 6\\%$. \\item Finally, we measure the colour of the nucleus to be $(V-R)=0.26\\pm0.09$, which is bluer than found for JFC nuclei but is affected by the activity. \\end{enumerate}"
    },
    "1005/1005.1675_arXiv.txt": {
        "abstract": "GRB~090417B was an unusually long burst with a $T_{90}$ duration of at least 2130~s and a multi-peaked light curve at energies of 15--150~keV.  It was optically dark and has been associated with a bright star-forming galaxy at a redshift of 0.345 that is broadly similar to the Milky Way.  This is one of the few cases where a host galaxy has been clearly identified for a dark gamma-ray burst and thus an ideal candidate for studying the origin of dark bursts.  We find that the dark nature of GRB~090417B cannot be explained by high redshift, incomplete observations, or unusual physics in the production of the afterglow.  Assuming the standard relativistic fireball model for the afterglow we find that the optical flux is at least 2.5~mag fainter than predicted by the $X$-ray flux.  The {\\sl Swift\\/}/XRT $X$-ray data are consistent with the afterglow being obscured by a dense, localized sheet of dust approximately 30--80~pc from the burst along the line of sight.  Our results suggest that this dust sheet imparts an extinction of $A_V \\ga 12$~mag, which is sufficient to explain the missing optical flux.  GRB~090417B is an example of a gamma-ray burst that is dark due to the localized dust structure in its host galaxy. ",
        "introduction": "}  The {\\sl Swift\\/} observatory \\citep{GCC2004} is a multi-instrument satellite mission that was designed to detect and rapidly localize gamma-ray bursts (GRBs).  The observatory contains three telescopes: the Burst Alert Telescope (BAT\\@; \\citet{BBC2005}), the $X$-Ray Telescope (XRT\\@; \\citet{BHN2005}), and the UltraViolet/Optical Telescope (UVOT\\@; \\citet{RKM2005}).  The BAT is used to identify GRBs and localize them to $\\sim 3\\arcmin$ in the energy range 15--150 keV. Once the BAT has localized a burst {\\sl Swift\\/} slews to point the XRT and the UVOT at the burst.  The XRT obtains rapid $X$-ray localizations to $\\lesssim 5\\arcsec$ in the energy range 0.2--10 keV while the UVOT obtains localizations to $\\approx 0\\farcs5$ then repeatedly cycles through a set of optical and ultraviolet filters. Over the past five years {\\sl Swift\\/} has detected GRBs at the rate of approximately 100 per year.  Almost all of these {\\sl Swift\\/}-detected GRBs have had $X$-ray afterglows, and optical or infrared afterglows have been detected for $\\approx 60$\\% of them. The remaining $\\approx$ 40\\% of the {\\sl Swift\\/} bursts have no reported optical or infrared detections.  In general between $\\approx 25$--40\\% of GRBs are dark GRBs \\citep{FJP2009}.  The gamma-ray burst GRB~090417B was detected by the BAT as an image trigger at 13:17:23 UT on 2009 Apr 17 \\citep{SBB2009}. {\\sl Swift\\/} immediately slewed to GRB~090417B and narrow-field observations began observing 387~s (XRT) and 378~s (UVOT) after the BAT trigger.  The BAT light curve showed a long period of emission starting 200~s before the trigger (see Figure~\\ref{FIGURE:batlc}).  Several ground-based follow-up observations were made, but no optical or infrared afterglow was detected.  \\citet{FMH2009} found a source with $R = 21.3$ inside the XRT error circle \\citep{SBB2009} and identified it as an SDSS galaxy \\citep{SWL2002} and the possible host galaxy of GRB~090417B.  \\citet{BF2009} measured a redshift of $z = 0.345$ for this galaxy based on the forbidden oxygen emission lines [\\ion{O}{2}](3727) and [\\ion{O}{3}](5006) as well as H$\\alpha$ emission.  Dark GRBs have been an enduring mystery in the lore of GRB studies. In the pre-{\\sl Swift\\/} era only about one quarter of localized gamma-ray bursts had an optical afterglow (OA) detected \\citep[\\eg,][]{FJG2001}.  It was sometimes assumed that this was due to the response times of ground-based optical telescopes which allowed most afterglows to fade before they could be observed.  {\\sl Swift\\/} however, with its ability to train the XRT and UVOT on a burst within $\\approx 100$~s of the initial detection, has been able to provide positions to within $\\approx 5\\arcsec$ for most bursts that have been detected by the BAT\\@.  The rapid distribution of these positions has allowed ground-based optical and infrared follow-up observations for many of these bursts, often within minutes of the BAT trigger.  Even with such rapid follow-up observations no optical or infrared afterglows have been found for $\\approx$ 40\\% of BAT-detected GRBs, thus demonstrating that observational constraints were not the cause to the dark burst problem \\citep{RSF2006}.  Several attempts have been made to quantify the dark burst problem by establishing operational definitions of dark bursts \\citep{JHF2004,RWK2005,HKG2009}.  Most recently \\citet{HKG2009} proposed that dark GRBs be defined as those that have $\\beta_{OX} < \\beta_X - 0.5$ measured 11 hours after the burst.  They define $\\beta_{OX}$ as the spectral index between the optical and $X$-ray regimes and $\\beta_X$ as the spectral index in the $X$-ray band.  The spectral index is defined by $f_\\nu(\\nu) \\propto \\nu^{-\\beta}$ where $f_\\nu(\\nu)$ is the flux density at a frequency $\\nu$.  This definition assumes that both the $X$-ray and optical afterglows are due to synchrotron radiation from a relativistic fireball.  The synchrotron model predicts that the spectral slope will either be the same in both regimes (if there is no cooling break between them) or will differ by $\\Delta\\beta = 0.5$ (if there is a cooling break between them).  A difference of $\\Delta\\beta > 0.5$ can not be explained in this model and thus requires that something be suppressing the flux at optical wavelengths, resulting in a dark burst \\citep{JHF2004}.  Several possibilities have been proposed to explain dark bursts. First, several studies have suggested that dark bursts may be due to extinction along the line of sight, either in the Milky Way or in the host galaxy \\citep[\\eg,][]{PFG2002,LFR2006,BFK2007,JRW2008,TLR2008}. \\citet{PCB2009} suggest that this is the case for most dark GRBs. Further, $X$-ray observations show that some bursts with no optical afterglows have systematically higher \\ion{H}{1} column densities than those with optical afterglows \\citep{FJP2009}.  This suggests that dark bursts may be suffering from higher extinction than optically-bright bursts.  However, there is significant overlap in the distribution of $N_{\\mathrm{H}}$ values between dark and bright bursts, so it is not clear if extinction is the sole parameter responsible for the darkness of these bursts.  \\citet{SMP2007,SPO2009} did panchromatic studies of several GRBs that were observed with {\\sl Swift\\/}'s XRT and UVOT ({\\ie}, optically-bright GRBs) in order to probe the extinction in their host galaxies.  They found that the combined $X$-ray/ultraviolet/optical data for most of these GRBs are best fit using a Small Magellanic Cloud (SMC) dust model \\citep{P1992}.  The extinctions in these galaxies were small, with $A_V < 1$ mag.  They also found a mean gas-to-dust ratio of $N_\\mathrm{H} = \\left(6.7 \\times 10^{21}\\right) A_V$, which is lower than, but similar to, the value found in the SMC ($N_\\mathrm{H} = \\left(15.4 \\times 10^{21}\\right) A_V$ using equation (4) and Table 2 of \\citet{P1992}).  \\citet{KKZ2006} also found that the the gas-to-dust ratio for a large sample of host galaxies of optically-bright GRBs is lower than the SMC value.  However, the gas-to-dust ratios in the host galaxies of dark GRBs tend to be more like that of the Milky Way than that of the SMC \\citep[\\eg,][]{PFG2002,BFK2007,JRW2008,EFH2009}.  A second possible origin for dark bursts is that they are GRBs located at high redshift, so the Lyman break (or the start of the Lyman-alpha forest) is observed redward of the optical band.  The most distant GRB observed to date is GRB~090423 at $z = 8.2$ \\citep{SDC2009,TFL2009}, so some of the dark bursts could simply be at high redshift.  However, \\citet{FJP2009} have found that less than approximately 19\\% of GRBs have $z > 7$, whereas between 25\\% and 42\\% of GRBs are dark. Therefore, it is unlikely that all of the {\\sl Swift\\/} dark burst lie at very high redshifts.  The affects of relativistic beaming of GRB emission at very high redshift is poorly understood due to a lack of data on high-redsift bursts.  We assume here that it does not affect our discussion in this paper.  Third, the hypothesis that afterglow radiation is synchrotron radiation may be wrong.  This could lead to a different spectrum from what is expected from the relativistic fireball model, in which case the above definition of a dark burst is invalid.  Long--soft GRBs \\citep{KMF1993} tend to be found in host galaxies that are small, irregular, and have high specific star-formation rates \\citep{CHG2004}.  \\citet{LDM2003} found a median infrared luminosity of $L \\approx 0.08 L^*$ for the host galaxies of GRBs with optical afterglows and conclude that most hosts are sub-luminous at optical and near-infrared wavelengths.  These GRB hosts tend to be small galaxies that have intrinsic luminosities similar to those of the dwarf galaxies in the Local Group.  However, some hosts are significantly larger, such as that of GRB~990705 \\citep{LDM2002}, which was an Sc spiral with $L \\approx 2L_\\star$.  To date no long--soft GRB has been found in an elliptical galaxy.  All of the hosts of long--soft GRBs have morphologies that are consistent with either exponential discs, or irregular structure \\citep{CVF2005,WBP2007}.  GRB hosts follow the size--luminosity relation \\citep[\\eg,][]{TFR2006} and extend it towards lower luminosities \\citep{WBP2007}.  In spite of their generally small size the host galaxies of long--soft GRBs tend to have high specific star-formation rates (SSFRs).  \\citet{SGL2009} found a median SSFR of 0.8 yr$^{-1}$, similar to what is seem in Lyman-break galaxies.  It is important to realize, however, that these properties have been determined using GRBs that have usually been localized to sub-arcsecond precisions.  In practice this usually means that the sample of GRB host galaxies is restricted to GRBs that have had optical afterglows.  The host galaxies of dark GRBs may have different properties.  \\citet{PCB2009} identify host galaxies for 14 dark GRBs and found that they have the same redshift distribution as the hosts of optically-bright GRBs.  They also found that these hosts do not appear to be significantly different in size, structure, or luminosity from the hosts of GRBs with optical afterglows.  They did find that a significant fraction of the hosts of dark bursts show evidence for high internal extinction ($A_V > 2$--5 mag).  However, many of their host associations are uncertain, with more than half of the galaxies having a probability of a chance association on the sky of $P_\\mathrm{ch} > 0.01$.  As we will show in this paper GRB~090417B is located $1\\farcs08$ from the centre of a bright galaxy with a probability of a chance association of only $P_\\mathrm{ch} \\approx 10^{-3}$.  This makes GRB~090417B one of the strongest associations of a dark GRB with a host galaxy.  In this paper we present space- and ground-based gamma-ray, $X$-ray, ultraviolet, and optical observations of GRB~090417B and its host galaxy.  We have adopted a standard ${\\Lambda}$CDM cosmology with a Hubble parameter of $H_0 = 71$ km s$^{-1}$ Mpc$^{-1}$, a matter density of $\\Omega_m = 0.27$, and a cosmological constant of $\\Omega_\\Lambda = 1 - \\Omega_m = 0.73$.  For this cosmology a redshift of $z = 0.345$ corresponds to a luminosity distance of 1.806 Gpc and a distance modulus of 41.29 mag.  One arcsecond corresponds to 6.51 comoving kpc, or 4.84 proper kpc.  The look back time is 3.80 Gyr. ",
        "conclusions": "}  We have shown that the host of the dark GRB~090417B is very likely to be SDSS~J135846.65+470104.5, an $L^*$ galaxy at $z = 0.345$.  This galaxy has an overall star-formation rate that is at least as great as that of the Milky Way, and an overall $V$-band extinction of $A_V \\approx 3.5$ mag, which is dustier than typical GRB host galaxies, but not greatly different from the Milky Way.  $X$-ray observations of GRB~090417B show the normal temporal and spectral behaviour seen in the $X$-ray afterglows of many {\\sl Swift\\/} GRBs.  The data up to approximately 70~ks after the BAT trigger are consistent with a relativistic fireball expanding into either a homogeneous ISM or a wind-stratified external medium.  The cooling frequency is below the $X$-ray band between 9.1~ks and 29.6~ks, and the $k$-corrected isotropic energy is $E_\\mathrm{iso} \\ga 6.3 \\times 10^{51}$ erg. The physics of this burst do not appear to be unusual, and thus are unlikely to explain the dark nature of the burst.  After approximately 70~ks the $X$-ray spectrum becomes significantly softer.  We find that this can be explained using the dust scattering model of \\citet{SD2007}.  The observed late-time spectral evolution can be produced by a sheet of dust approximately 30--80 pc from the burst along the line of sight.  The model predicts an extinction of $A_V \\approx 15$--40~mag along the line of sight. This is consistent with the \\ion{H}{1} column density measured from the $X$-ray spectrum.  Therefore, we conclude that GRB~090417B is probably dark because of a dense, localized layer of dust along the line of sight between us and the afterglow."
    },
    "1005/1005.4785_arXiv.txt": {
        "abstract": "{The origin of hot subdwarf B stars (sdBs) is still unclear. About half of the known sdBs are in close binary systems for which common envelope ejection is the most likely formation channel. Little is known about this dynamic phase of binary evolution. Since most of the known sdB systems are single-lined spectroscopic binaries, it is difficult to derive masses and unravel the companions' nature, which is the aim of this paper.\\\\ Due to the tidal influence of the companion in close binary systems, the rotation of the primary becomes synchronised to its orbital motion. In this case it is possible to constrain the mass of the companion, if the primary mass, its projected rotational velocity as well as its surface gravity are known. For the first time we measured the projected rotational velocities of a large sdB binary sample from high resolution spectra. We analysed a sample of 51 sdB stars in close binaries, 40 of which have known orbital parameters comprising half of all such systems known today.\\\\ Synchronisation in sdB binaries is discussed both from the theoretical and the observational point of view. The masses and the nature of the unseen companions could be constrained in 31 cases. We found orbital synchronisation most likely to be established in binaries with orbital periods shorter than $1.2\\,{\\rm d}$. Only in five cases it was impossible to decide whether the sdB's companion is a white dwarf or an M dwarf. The companions to seven sdBs could be clearly identified as late M stars. One binary may have a brown dwarf companion. The unseen companions of nine sdBs are white dwarfs with typical masses. The mass of one white dwarf companion is very low. In eight cases (including the well known system KPD1930$+$2752) the companion mass exceeds $0.9\\,M_{\\rm \\odot}$, four of which even exceed the Chandrasekhar limit indicating that they may be neutron stars. Even stellar mass black holes are possible for the most massive companions. The distribution of the inclinations of the systems with low mass companions appears to be consistent with expectations, whereas a lack of high inclinations becomes obvious for the massive systems. We show that the formation of such systems can be explained with common envelope evolution and present an appropriate formation channel including two phases of unstable mass transfer and one supernova explosion. The sample also contains a candidate post-RGB star, which rotates fast despite its long orbital period. The post-RGB stars are expected to spin-up caused by their ongoing contraction. The age of the sdB is another important factor. If the EHB star is too young, the synchronisation process might not be finished yet. Estimating the ages of the target stars from their positions on the EHB band, we found PG\\,2345$+$318, which is known not to be synchronised, to lie near the zero-age extreme horizontal branch as are the massive candidates PG\\,1232$-$136, PG\\,1432$+$159 and PG\\,1101$+$249. These star may possibly be too young to have reached synchronisation.\\\\ The derived large fraction of putative massive sdB binary systems in low inclination orbits is inconsistent with theoretical predictions. Even if we dismiss three candidates because they may be too young and assume that the other sdB primaries are of low mass, PG\\,1743$+$477 and, in particular, HE\\,0532$-$4503 remain as candidates whose companions may have masses close to or above the Chandrasekhar limit. X-ray observations and accurate photometry are suggested to clarify their nature. As high inclination systems must also exist, an appropriate survey has already been launched to find such binaries. ",
        "introduction": "}  Subluminous B stars (sdBs) show similar colours and spectral characteristics as main sequence stars of spectral type B, but are much less luminous. Compared to main sequence B stars the hydrogen Balmer lines in the spectra of sdBs are stronger while the helium lines are much weaker (if present at all) for the colour. The strong line broadening and the early confluence of the Balmer series is caused by the high surface gravities ($\\log\\,g\\simeq5.0-6.0$) of these compact stars ($R_{\\rm sdB}\\simeq0.1-0.3\\,R_{\\rm \\odot}$). Subluminous B stars are considered to be helium core burning stars with very thin hydrogen envelopes and masses of about half a solar mass (Heber \\cite{heber1}) located at the extreme end of the horizontal branch (EHB).  Subdwarf B stars are found in all Galactic stellar populations and are sufficiently common to account for the UV-upturn of early-type galaxies. Understanding the origin of the UV-upturn phenomenon hence has to await a proper understanding of the origin of the sdB stars themselves.  The discovery of short-period multi-periodic pulsations in some sdBs provided an excellent opportunity to probe the interiors of these stars using the tools of asteroseismology. They were theoretically predicted by Charpinet et al. (\\cite{charpinet96}) at around the same time as they were observed by Kilkenny et al. (\\cite{kilkenny}). They are characterised by low-amplitude, multi-periodic, short-period ($80-600\\,{\\rm s}$) light variations that are due to pressure ($p$)-mode oscillations. A second family of pulsating sdB stars was discovered by Green et al. (\\cite{green2}), again showing low-amplitude, multi-periodic pulsations, but periods are longer ($2000-9000\\,{\\rm s}$) and are identified as gravity ($g$) modes. An important recent achievement of sdB asteroseismology is the determination of the most fundamental parameter of a star, i.e. its mass (for a review see Fontaine et al. \\cite{fontaine}).  The origin of EHB stars, however, is wrapped in mystery (see Heber \\cite{heber09} for a review). The problem is how some kind of mass loss mechanism in the progenitor manages to remove all but a tiny fraction of the hydrogen envelope at about the same time as the helium core has attained the mass ($\\sim0.5$ M$_\\odot$) required for the helium flash. This requires enhanced mass loss, e.g. due to helium mixing driven by internal rotation (Sweigart \\cite{sweigart97}) or at the helium flash itself.  Mengel et al. (\\cite{mengel76}) demonstrated that the required strong mass loss can occur in a close-binary system. The progenitor of the sdB star has to fill its Roche lobe near the tip of the red-giant branch (RGB) to lose most of its hydrogen-rich envelope. The merger of binary white dwarfs was investigated by Webbink (\\cite{webbink}) who showed that an EHB star can form when two helium core white dwarfs merge and the product is sufficiently massive to ignite helium.  Interest in the binary scenario was revived, when Maxted et al. (\\cite{maxted2}) determined a very high fraction of radial velocity variable sdB stars, indicating that about two thirds of the sdB stars in the field are in close binaries with periods of less than 30 days (see also Morales-Rueda et al. \\cite{morales}; Napiwotzki et al. \\cite{napiwotzki8}). The companions, as far as their nature could be clarified, are mostly M dwarfs or white dwarfs. If the white dwarf companion is sufficiently massive, the merger of the binary system might exceed the Chandrasekhar mass and explode as a type Ia supernova. Indeed,  Bill\\`{e}res et al. (\\cite{billeres00}) and Maxted et al. (\\cite{maxted}) discovered KPD~1930$+$2752, a system that might qualify as a SN Ia supernova progenitor (see also Geier et al. \\cite{geier}).  These discoveries triggered new theoretical evolutionary calculations in the context of binary population-synthesis to identify the importance of various channels of close-binary evolution (Han et al. \\cite{han1,han2}), i.e. two phases of common-envelope ejection, stable Roche-lobe overflow and white dwarf merger.  \\subsection{Outline of the paper}  The purpose of this paper is to clarify the nature of the unseen companions for 40 short-period sdB binaries, which comprises about half of the sdB stars in single-lined close binary systems with known periods and radial velocity amplitudes. We assumed tidally locked rotation and made use of the sdBs' gravities and projected rotational velocities.  The paper is structured in two parts. After a short review on close binary sdB stars (Sect.~\\ref{sec:binaries}), part I (Sects.~\\ref{sec:obs} to \\ref{sec:form}) describes the analysis of the sample. Besides constraining the mass of the companions and unravelling the nature of most companions as M dwarfs or typical white dwarfs, it reports the discovery of a population of eight unseen compact companions with masses exceeding $0.9\\,M_{\\rm \\odot}$ (in addition to KPD~1930$+$2752), some of which even exceed the Chandrasekhar limit. Accordingly, the latter should be neutron stars (NS) or black holes (BH). Even if they were massive white dwarfs, it would be surprising to find such a large fraction, as massive white dwarfs are rare. As no binary system containing an sdB plus a NS/BH is known today, we investigate potential formation scenarios in Sect.~\\ref{sec:form} and find it indeed possible to create such systems through two phases of common envelope evolution.  Our results rest on the assumption of tidally locked rotation. Therefore, part II of the paper (Sects.~\\ref{sec:tidal} and \\ref{sec:challenge}) deals with the synchronisation time scales of sdB stars in close binaries both from a theoretical point of view and from the perspective of empirical constraints. The general result is that systems with periods shorter than $1.2\\,{\\rm d}$ should be synchronised. Empirical evidence is available that systems with periods below $0.6\\,{\\rm d}$ are synchronised as is indeed the case for the systems with massive companions.  Although selection effects would favour detection of highly inclined systems, no such system was found among those binaries with massive companions in our sample. This calls for a careful inspection of alternative explanations (Sect.~\\ref{sec:challenge}). There are two aspects to be discussed. First, the sdB may not burn helium at all and, thus, is spun up due to ongoing contraction. Alternatively the actual evolutionary age of individual stars may be smaller than appreciated, i.e. the EHB star may just have formed only recently and the systems would therefore not be synchronised. In Sect.~\\ref{sec:summary} we summarise and discuss the results. ",
        "conclusions": "\\label{sec:summary}  We have analysed a sample of 51 sdB stars in close single-lined binary systems. This included 40 systems for which the orbital parameters have been determined previously. The subsample comprises half of all systems known so far. From high resolution spectra taken with different instruments the projected rotational velocities of these stars have been derived to an unprecedented precision. Accurate measurements of the surface gravities have mostly been taken from literature. Assuming orbital synchronisation and an sdB mass distribution as suggested by binary population synthesis models as well as by asteroseismology, the masses and the nature of the unseen companions could be constrained in 31 cases. Only in five cases we were unable to classify unambiguously. These companions may either be low mass main-sequence stars or white dwarfs. The companions to seven sdBs could be clearly identified as late M stars. One binary may have a brown dwarf companion. The unseen companions of nine sdBs are white dwarfs with typical masses, one WD companion has a very low mass.  In eight cases (including the well known system KPD1930$+$2752) the companion mass exceeds $0.9\\,M_{\\rm \\odot}$. Four of the companions even exceed the Chandrasekhar limit indicating that they may be neutron stars; even a stellar mass black hole is possible for the most massive companions.  The basic assumption of orbital synchronisation in close sdB binaries has been discussed in detail. Our analysis method yielded consistent results for binaries up to an orbital period of $\\simeq1.2\\,{\\rm d}$. Theoretical timescales for synchronisation were calculated using two different approaches. The theory of Zahn (\\cite{zahn2}) was found to be too inefficient while that of Tassoul \\& Tassoul (\\cite{tassoul}) predicts too short timescales. The predictions from both theories are strongly discrepant, calling for empirical constraints.  Independent observational evidence for synchronisation in sdB binaries comes from light curve analyses of eclipsing, ellipsoidal deformed, and pulsating sdBs. Due to this evidence sdB binaries with periods shorter than $\\simeq0.6\\,{\\rm d}$ should be synchronised. This includes all of the putative massive systems.  Hence, an evolutionary model for the origin of sdB stars with neutron star or black hole companions was devised indicating that common envelope evolution is indeed capable of producing such systems, though at a lower rate than observed. An appropriate formation channel includes two phases of unstable mass transfer and one supernova explosion.  The distribution of the inclinations of the systems of normal mass appears to be consistent with expectations, whereas a lack of high inclinations became obvious for the massive systems.  There is one star in the sample which rotates fast despite its rather long orbital period. This as well as its position far from the EHB band hints at a post-RGB nature. The post-RGB stars are expected to be spun-up due to their ongoing contraction.  The larger number of putative massive companions in low inclination systems is puzzling. Therefore, we investigated alternative interpretations. The fraction of massive unseen companions can only be lowered, if the sdBs themselves have masses much lower than the anticipated range of $0.43 - 0.47\\,M_{\\rm \\odot}$ for EHB stars. Evolutionary calculations showed that EHB stars with masses as low as $0.30\\,M_{\\rm \\odot}$ can be formed if helium ignites under non-degenerate conditions but should be very rare. Assuming such low sdB masses, only one unseen companion remains more massive than the Chandrasekhar limit. This fraction of $3\\%$ is roughly consistent with theoretical predictions. Whether the sdB mass is small or not can be checked directly as soon as accurate parallaxes of these relatively bright stars will become available through the GAIA mission.  The putative massive sdB systems might not be synchronised if their age is much less than anticipated. That this can happen is witnessed by PG\\,2345$+$318, a short-period sdB binary in our sample, that we would have classified as a low-inclination massive system as well, if it were not proven by eclipses to be highly inclined. Hence the system is not synchronised despite of its short period ($0.24\\,{\\rm d}$). Due to a degeneracy between mass and age, it is difficult to estimate the sdB's age without knowing its mass. Adopting the canonical mass, we nevertheless estimated the stars' ages from their position in the EHB band. Indeed, PG\\,2345$+$318, is located right on the zero-age EHB as are the massive candidates PG\\,1232$-$136, PG\\,1432$+$159 and PG\\,1101$+$249. These stars may possibly be too young to have reached synchronisation. Hence the companion masses we derived would be spurious. However, there is no indication that the other massive systems could be young.  Even if we dismiss three candidates because they may be too young and assume that the others are of low mass, PG\\,1743$+$477 and, in particular, HE\\,0532$-$4503 remain as massive candidates whose companions have masses close to or above the Chandrasekhar mass.  Different approaches may be chosen to directly verify the presence of neutron star or black hole companions in our candidate systems. None of the sdBs in our target systems fills its Roche lobe. No mass transfer by Roche lobe overflow to the unseen companion can occur and therefore no X-ray emission is expected. The ROSAT all-sky survey catalogue (RASS, Voges et al. \\cite{voges}) has been checked and, indeed, no sources have been detected at the positions of any candidate sdB+NS/BH systems. The detection limit of this survey reaches down to about $10^{-13}\\,{\\rm erg\\,cm^{-2}s^{-1}}$. However, sdB stars are expected to have weak winds. Hence accretion from the sdB wind might result in faint X-ray emission. This occurs in the bright sdO+WD system HD\\,49798 (Mereghetti et al., \\cite{mereghetti}). Although stellar wind mass loss rates in sdBs are predicted to be small ($<10^{-12}\\,M_{\\rm \\odot}{\\rm yr^{-1}}$, e.g. Vink \\& Cassisi \\cite{vink}; Unglaub \\cite{unglaub3}), they may be sufficient to cause detectable X-ray flux powered by wind accretion.  X-ray telescopes like Chandra or XMM-Newton may be sensitive enough to detect such weak sources. Pulsar signatures of rapidly spinning neutron star companions may be detectable with radio telescopes.  Tidal forces by the companion cause an ellipsoidal deformation of the primary in close binary systems. This deformation appears as a variation of light at half the orbital period. Two very close subdwarf binaries with orbital periods of $\\simeq\\,2\\,{\\rm h}$ and high orbital inclination show light variations of about $1\\%$, which can be detected from the ground. Performing binary light curve synthesis it was possible to derive the masses of the binary components (Orosz \\& Wade \\cite{orosz}; Geier et al. \\cite{geier}). Signatures of ellipsoidal deformation in the light curves of binaries with longer orbital periods and lower inclination are much weaker ($\\simeq0.01\\%$, Drechsel priv. comm.; Napiwotzki et al. in prep.) and therefore not detectable from the ground. The existence of such very shallow variations has been proven for the subdwarf binary PG\\,0101$+$039 with an orbital period of $13.7\\,{\\rm h}$ using a light curve of almost $17\\,{\\rm d}$ days duration taken with the MOST satellite. The ellipsoidal variation was found to be $0.025\\%$ (Geier et al. \\cite{geier2}). The full potential of high precision photometry for the analysis of sdB binaries has most recently been demonstrated by Bloemen et al. (\\cite{bloemen}), who analysed a Kepler light curve of the eclipsing sdB+WD binary KPD\\,1946$+$4340. High precision light curves of the best candidates in our sample should be measured with HST. The nature of their unseen companion could then be clarified.  Most of the candidate massive systems have low orbital inclination. High inclination systems must exist as well. In this case a determination of the orbital parameters is sufficient to put a lower limit to the companion mass by calculating the binary mass function. If this lower limit exceeds the Chandrasekhar mass and no sign of a companion is visible in the spectra, the existence of a massive compact companion is proven without making any additional assumptions. The Hyper-MUCHFUSS project (Hypervelocity stars or Massive Unseen Companions to Hot Faint Underluminous Stars from SDSS, Geier et al. in prep.) was launched in 2007. One of the aims of this project is to search for sdB binaries with massive compact companions at high inclinations in a sample of stars selected from the SDSS data base."
    },
    "1005/1005.4050_arXiv.txt": {
        "abstract": "Radial velocity measurements of stellar reflex motion have revealed many extrasolar planets, but gaps in the observations produce aliases, spurious frequencies that are frequently confused with the planets' orbital frequencies. In the case of Gl~581~d, the distinction between an alias and the true frequency was the distinction between a frozen, dead planet and a planet possibly hospitable to life \\citep{2007U,2009M}. To improve the characterization of planetary systems, we describe how aliases originate and present a new approach for distinguishing between orbital frequencies and their aliases. Our approach harnesses features in the spectral window function to compare the amplitude and phase of predicted aliases with peaks present in the data. We apply it to confirm prior alias distinctions for the planets GJ~876~d and HD~75898~b. We find that the true periods of Gl~581~d and HD~73526~b/c remain ambiguous. We revise the periods of HD~156668~b and 55~Cnc~e, which were afflicted by daily aliases.  For HD~156668~b, the correct period is 1.2699 days and minimum mass is $(3.1 \\pm 0.4)$~$M_\\oplus$. For 55~Cnc~e, the correct period is 0.7365 days -- the shortest of any known planet -- and minimum mass is $(8.3 \\pm 0.3)$~$M_{\\oplus}$. This revision produces a significantly improved 5-planet Keplerian fit for 55 Cnc, and a self-consistent dynamical fit describes the data just as well.  As radial velocity techniques push to ever-smaller planets, often found in systems of multiple planets, distinguishing true periods from aliases will become increasingly important. ",
        "introduction": "\\label{sec:intro}  In the past two decades, over 400 extrasolar planets have been discovered, including more than 300 detected by radial velocity measurements. The entire architecture of a planetary system is encoded in the wobbles of its host star. In frequency space, the star's radial velocity variations are decomposed into the frequencies associated with each planet's gravitational interactions. One obstacle in correctly attributing these frequencies to planets are the spurious alias frequencies in the periodogram of the star's radial velocity measurements, caused by the discrete time sampling of the observations. Convolved with the orbital frequencies of alien worlds are Earth's own rotational and orbital frequencies, which dictate when the host star is visible at night, and -- for many data sets -- the synodic lunar frequency, which impacts the allocation of telescope time.  Distinguishing aliases from physical frequencies is a common problem, yet making the correct distinction is crucial for characterizing extrasolar planets. For example, \\citet{2007U} announced a super-Earth orbiting the M star Gl~581 with period 83 days, beyond the cold edge of the habitable zone. After more than doubling the number of observations, they determined that the planet's period was actually 67 days, well within the habitable zone, and that the 83 day period was an alias \\citep{2009M}. The distinction between an alias and physical frequency was the distinction between a frozen, dead planet and a planet  possibly hospitable to life. For reasons we will describe below, planets with periods of one to several months -- in or near the habitable zone of M stars -- will typically have aliases with periods within about 30 days of their own orbital period. As more planets are discovered orbiting M stars, astronomers will be struggling to distinguish which of two close frequencies, one of which places the planet in the habitable zone, corresponds to a planet's orbital frequency. In general, planets with periods between a few months and a few years often have confusing aliases caused by convolution with Earth's orbital period, while planets with periods near a day, such as the super-Earth GJ~876~d \\citep{2005R}, have confusing aliases caused by convolution with Earth's rotational period. Automatic de-aliasing algorithms, such as CLEAN \\citep{1987R}, have been applied to particularly complicated radial velocity periodograms with some success \\citep{2009Q}, yet, while they are good for cleaning up a periodogram, they should not be relied on for distinguishing between an alias and a physical frequency.  Aliases also pose a challenge for observing variable stars and period-searching algorithms have been designed to not fall prey to them (see for example \\citealt{2008P, 2007Reegen,2010R}).  Therefore, to enhance detection and characterization of planets, we have developed an approach to identify aliases by harnessing features of the ``spectral window function,\" the Fourier transform of the observation times. Consider the star's motion as a signal that passes through a system, the time sampling window. Because of noise and loss of information, we can never perfectly reconstruct the signal. But we know everything there is to know about the system: for a sinusoid of a given amplitude, frequency, and phase, peaks in the window function cause aliases with calculable amplitudes and phases \\citep{1975D,1976D}. The several time sampling frequencies -- sidereal year, sidereal day, solar day, and synodic month -- complicate the radial velocity periodogram yet allow us to break the degeneracy between alias and physical frequency that would exist for evenly-sampled data.  In the following section, we describe the origin and characteristics of aliases, supply the details of our approach for confirming that a particular frequency is not an alias, and clarify previous misconceptions about aliases. In the third section, we apply our approach to confirm periods for the planets GJ~876~d and HD~75898~b. We find that the orbital period for Gl~581~d and for the planets of HD~73526 cannot be definitively determined due to noise. We discover that the reported orbital period for HD~156668~b, 4.6455 days, is an alias of the true period,  1.2699 days. Finally, we analyze the 5-planet system 55 Cnc. We find that the period of 2.817 days reported in the literature for planet e \\citep{2004M, 2008F} is actually a daily alias of its true period of 0.737 days. We conclude by summarizing the approach we have developed, considering the implications of a new period for 55~Cnc~e, and suggesting observational strategies for mitigating aliases. ",
        "conclusions": "\\subsection{Summary of Approach}  Aliases result from a convolution between a true physical frequency and the spectral window function, which is created by gaps in the data set due to observational constraints. Our method harnesses features in the window function to distinguish aliases from true frequencies. For a given frequency $f$ and window function peak $f_s$, aliases will occur at $|f \\pm f_s |$, where $f_s$ is a feature in the window function. In the ranges where we expect major aliases to occur, we compare the phase and amplitude of aliases predicted by a sinusoid of the candidate frequency sampled to the data, with other known planets subtracted off beforehand. We judge whether the ``pattern\" of the predicted aliases matches the data: for example, yearly aliases appear as sidebands of the candidate frequency while daily aliases often appear as a doublet caused by the sidereal and solar day. If all the aliases match in amplitude, phase, and pattern, we can be confident that we have found the true orbital period. If there are discrepancies and the aliases of none of the candidate frequencies match the data, we know that noise prevents us from definitively determining the true period and that follow-up observations are necessary. Misunderstandings about aliases have previously led to incorrect identification of planet's orbital periods, a key parameter in defining the planet's properties, as well as the dynamical behavior of the planets in the system. We have corrected common misconceptions, including that aliases always appear near the frequency of peaks in the window function, that any frequency above 1 cycle/day is necessarily an alias, and that aliases will appear if the data are scrambled or if the true frequency is subtracted out.  \\subsection{Summary of Results}  For two systems, we confirmed previous distinctions between alias and true frequency. The period of GJ~876~d is indeed 1.94 days, not 2.05 days. The period of HD~75898~b is indeed 400 days and the periodogram peak at 200 days is indeed an alias, not a second planet or eccentricity harmonic, the alternative explanations proposed by \\citet{2007R}.  For two other systems, we determined that the data are too noisy to allow us to definitely distinguish between alias or true frequency. According to our analysis, it remains unclear whether the period of Gl~581~d is 67 days or 83 days; even a period of ~1 day cannot be ruled out. It also remains unclear whether HD~73526 contains two planets with orbital periods 187.5 and 376.9 days, locked in a 2:1 resonance, or whether one of the periods is actually 127 days. Further observations of these systems are required, preferably at times that reduce the aliasing.  For a final pair of systems, we determined the reported orbital period was incorrect, due to mistaking a daily alias for the true frequency. According to our analysis, the orbital period of HD~156668~b is actually 1.2699 days, not 4.6455 days. The orbital period 55~Cnc~e is 0.7365 days, not 2.817 days. The standard, general-purpose software \\emph{SigSpec} mentioned in the introduction \\citep{2007Reegen,2010R} agrees with our orbital period distinctions (we used the parameters: depth=2, par = 0.2 and par = 0.5, and a frequency upper limit of 2 day${-1}$).  \\subsection{Implications for 55~Cnc~e}  What are the implications of an updated period for the innermost planet of 55 Cnc?  First, it dramatically lowers the effective noise when determining the parameters of the planetary system.  \\citet{2008F} reported independent Keplerian fits with rms of 6.74 m/s, and a self-consistent dynamical fit with rms of 7.712 m/s. Our Keplerian fit achieves rms of 5.91 m/s, and our self-consistent coplanar dynamical fit achieves rms of 5.96 m/s. By adjusting the inclination of the system relative to our line-of-sight and the planets' mutual inclinations, an even better self-consistent might be possible. Therefore perturbations might be directly detected via a lower rms when interactions among the planets are included, and the architecture of the system further constrained.  We have just begun exploring this avenue.  Second, 55~Cnc~e itself can now be searched for transits at the new period, with high a priori probability of $\\sim25$\\%.  Given the period and phase of the radial-velocity signal, we report predicted transit epochs in Tables \\ref{tab:55cnc_newepoch} and \\ref{tab:55cnc_newepoch_ee0}.  The predictions differ because the latter assumes zero eccentricity, and the formally significant value of $e_{\\rm e}$ matters.  Nevertheless, folding the systematic uncertainty related to eccentricity into the predicted transit time, we still can predict transit times good to $\\sigma_T\\simeq1$~hour in 2010.  This search can be accomplished simply by folding the photometric data reported by \\cite{2008F} at the new ephemeris.  Gregory Henry (priv. comm.) has made such a search, finds no positive signal, and constrains putative transits in the period range $0.7-0.8$~days to a depth $<0.7$~mmag, or $>2.6$~$R_\\oplus$.  Earth-composition models of super-Earths predict a radius $\\sim1.9$~$R_\\oplus$ \\citep{2006V}, so a search at higher precision is certainly worthwhile.  Third, even apart from a transit, this super-Earth must be very hot, as it is very close-in to a solar-type star. Following \\citet{2009L}, we find that the substellar point could be up to 2750 K, if the insolation is absorbed then reradiated locally.  We would naively expect that the enormous radiation this planet takes in would evaporate any atmosphere (e.g., \\citealt{2010J}).  Moreover, the host star is also very bright as seen from Earth.  Therefore it might be useful to look for its phase curve with Spitzer, to detect or rule out an atmosphere \\citep{2009SD}.  Another attractive possibility is probing a magma ocean, which may exist because of the irradiation \\citep{2009G,2010G}, but this may require transit measurements.  Fourth, the presence of the other 4 planets surely injects a non-zero eccentricity into this tidally-dissipating planet.  Its expected value remains to be calculated, but will likely be on the order of $10^{-4}$.  This forced eccentricity could stimulate considerable geologic activity --- it might be a ``super-Io'' \\citep{2010B}.  \\subsection{Observational Strategies for Mitigating Aliases}  Can aliases be prevented or mitigated by the choice of observation times? Constraints on when the star is visible at night necessarily result in gaps in the data that cause aliases. However, we encourage observers to engage in ``window carpentry\" \\citep{1982S} by observing the star during the greatest span of the sidereal and solar day possible, not just when the star transits the meridian. Unfortunately, observing stars as they rise and set poses a challenge for observers, who minimize slew time\\footnote{If the slew time exceeds the read-out time, fewer observations may be gathered per night.  However, the wise spacing of observation times can more than make up for this through disambiguation of alias frequencies using fewer data points.} and thus maximize the number of stars observed per night by observing at the meridian for the majority of the night. This observing strategy (Fig.~\\ref{fig:mw1}b) results in strong daily and yearly aliases (ex. Fig.~\\ref{fig:wingl581}). Another strategy is to start in the west and gradually move east over the course of the night (Fig.~\\ref{fig:mw1}c), observing as much of the sky as possible. This strategy reduces yearly aliases but  sidereal daily aliases remain strong (ex. Fig.~\\ref{fig:winhd156668} and  \\ref{fig:win55cncfischer}). To reduce sidereal daily aliases, we recommend the following procedure.  Start the telescope somewhere west of the meridian (randomized from night to night) and move east to cover half the sky over the course of half the night. Then make one large slew to the place the telescope started and re-observe the same portion of the sky. Some stars will gain the advantage of being observed twice in one night.  Moreover, when the data are folded at the mean sampling period, they still show some variety in phase of observation, which is needed to reduce window function peaks and de-alias candidate periods. However, another consideration is that at higher air mass, both the extinction is greater and the seeing is worse. The increased atmospheric attenuation means a longer integration time is required, reducing the number of stars that can be observed, while the seeing increases the measurement errors. For a particular set of stars, observers can work out a slew pattern that will maximize the number of stars observed while minimizing aliasing. \\citet{2006S} present a clever method for determining the optimal sampling when period searching using satellite telescopes or a longitude-distributed network that can observe continuously. Unfortunately this strategy is impractical to implement using a single telescope on the ground. \\citet{2008Ford} presents useful adaptive scheduling algorithms for observing multiple targets that can be parametrized to reduce aliasing.   We suggest taking advantage of any unusual time windows: for example, the rare granted dark time or time at the beginning or end of another observer's night. Observers focusing on a large group of stars can determine which star would most benefit from this unusual time by calculating the window function with the new observation times added or, in the case of a planet with two candidate periods, determining for which system the observation times would best distinguish between two candidate orbits. We also suggest that it would be beneficial to observe stars using telescopes in two or more locations at different latitude and longitudes (ex. Fig.~\\ref{fig:win55cnccomb}).  At the stage of data analysis, we encourage the use of our method to distinguish true frequencies from aliases, crucial for the correct characterization of the planet. As astronomers push to observing lower mass planets and modeling planets near the noise limit, they cannot assume that the highest peak in the periodogram -- or even the best Keplerian fit -- corresponds to the true orbital period. Only by harnessing features in the window function to compare the amplitude, phase, and pattern of an assortment of predicted aliases to the data can we distinguish the planet's true orbital frequency -- or determine that more observations are needed.  \\subsection{Conclusion}  Knowing a planet's correct orbital period is essential for accurately characterizing it. By Kepler's law, the planet's distance from the star increases as its orbital period increases. Therefore the planet's orbital period sets its temperature: too hot, too cold, or just right for life. The planet's inferred mass, as calculated from the radial velocity amplitude, increases as the period decreases -- a closer planet needs less mass to exert a given force on the star -- so a difference in orbital period may be the difference between an Earth analog and a super-Earth. In the case of multi-planet systems, the spacing of the planets determines their mutual interactions: therefore a difference in orbital period may be the difference between a precariously placed planet and one locked deep in a stabilizing resonance. The signal of a planet's eccentricity is contained in the harmonics of the planet's orbital period: therefore a difference in orbital period may be the difference between a planet that formed in situ and a planet violently scattered, a calm planet that has long been tidally circularized or a planet erupting with volcanoes due to tidal dissipation. But periods that correspond to totally different worlds are only subtly distinguishable in the radial velocity signal. Such are the machinations of aliases.  Through our method, astronomers can confirm a planet's orbital period or determine that noise prevents a definitive distinction. In the latter case, follow-up observations taken according to the suggestions above should eventually allow the true period to be determined. Ironically, Earth's own rotational and orbital period make it challenging to uncover the orbital period of other worlds, particularly Earth analogs. But by better understanding of digital signal processing, we can mitigate the deleterious effects of the inevitable sunrise and starset."
    },
    "1005/1005.0862_arXiv.txt": {
        "abstract": "We present optical spectra and high-resolution multi-wavelength radio observations of the compact steep-spectrum radio source MRC~B1221$-$423 ($z=0.1706$).  MRC~B1221$-$423 is a very young ($\\sim 10^5$~yr), powerful radio source which is undergoing a tidal interaction with a companion galaxy.  We find strong evidence of interaction between the AGN and its environment. The radio morphology is highly distorted, showing a dramatic interaction between the radio jet and the host galaxy, with the jet being turned almost back on itself.  \\hi\\ observations show strong absorption against the nucleus at an infall velocity of $\\sim 250\\,$\\kms\\ compared to the stellar velocity, as well as a second, broader component which may represent gas falling into the nucleus.  Optical spectra show that star formation is taking place across the whole system. Broad optical emission lines in the nucleus show evidence of outflow. Our observations confirm that MRC~B1221$-$423 is a young radio source in a gas-rich nuclear environment, and that there was a time delay of a few $\\times 100$~Myr between the onset of star formation and the triggering of the AGN. ",
        "introduction": "\\label{sec:intro}  There is growing evidence of strong links between the triggering of powerful radio sources and gravitational interactions between galaxies.  Many radio galaxies show signatures of tidal interactions, such as tails, bridges, shells, and double nuclei \\citep[e.g.][]{hsb+86,bh92}.  These tidal encounters can lead to large increases in the amount of material being fed to the central black hole, thereby triggering the strong radio emission.  The same tidal event that dumps gas in the vicinity of the nuclear black hole can produce a burst of star formation, and hence a population of young stars.  Further, the mechanical energy associated with relativistic radio jets can have a profound effect on the formation and evolution of galaxies, by heating the interstellar gas to the point where stars can no longer form. This \\emph{radio-mode feedback} has been suggested as the mechanism which suppresses star formation in massive galaxies; without such a feedback mechanism, galaxy evolution models do not match the observed properties of galaxies \\citep{csw+06}.  So both starbursts and the growth of the nuclear black hole are intimately associated with the merger history of the galaxy.  However, key questions remain unanswered. If mergers are responsible for the formation of both AGN and starbursts, why do we sometimes see one without the other? When we do see both, we usually see a young radio source together with an older starburst, suggesting that there is a delay between the starburst and the onset of AGN activity.  What is the reason for this delay?  The compact steep-spectrum (CSS) sources are ideal laboratories for testing these connections.  CSS sources are powerful but compact radio sources, with sizes 1--20~kpc, so that the radio source lies well within its host galaxy. Their radio spectra are steep, $\\alpha < -0.5$\\ ($S_\\nu \\propto \\nu^\\alpha$), peaking below 500~MHz.  They have the same distribution in redshift as the rest of the powerful radio source population \\citep[see][for a review]{ode98}.  The weight of evidence now is that CSS sources are small because they are young; if fuelling is sustained they are expected to evolve into classical large double-lobed radio sources \\citep{ode98}. All CSS sources accessible to detailed optical investigations show evidence for recent interactions and/or mergers, and evidence for young stellar populations and galaxy-wide starbursts have been found in many of these objects \\citep[see][for a review]{hol09}.  MRC~B1221$-$423 is a member of a sample of southern radio sources similar in power to 3C sources \\citep{bh06}, with a radio power of $P_{\\mathrm{1.4 GHz}} = 1.8 \\times 10^{26}\\;{\\rm W\\,Hz}^{-1}$. The optical emission of the host galaxy is distorted, showing clear signatures of interaction, including tidal tails and shells \\citep{shp03}.  Multicolour images show knotty blue star-forming regions, while the likely merging galaxy to the south is also significantly bluer than the host galaxy (Johnston et al. 2005; see also Figure~\\ref{fig:slitpos})\\nocite{jhcs05}. With a redshift of $z=0.1706$\\ \\citep{scrw93}, MRC~B1221$-$423 is one of the nearest CSS sources; in the 3C sample, we need to go to $z=0.3$\\ to find the first compact source as powerful as MRC~B1221$-$423.  The radio source has double lobes which lie well within the envelope of the host galaxy \\citep{shp03}.  The age of the radio source, estimated from the power of the radio jets and total stored energy in the synchrotron plasma of the lobes \\citep{rs91}, is only $\\sim 10^5$~yr. As a young, nearby, bright CSS source, MRC~B1221$-$423 is therefore an ideal candidate for testing the relationship between the AGN and its environment.  The goal of this study is to understand the relationship between the central AGN, which is the observational signpost of a growing black hole, and the environment of the galaxy. MRC~B1221$-$423 provides a unique opportunity to measure both the age of the radio source and the star formation history of the galaxy, both of which are connected to the merger history of the galaxy.  In a previous paper (\\citealt{jhcs05}, henceforth Paper I) we presented an analysis of multi-colour \\emph{BVRIK} images of MRC~B1221$-$423, which we used to analyse the change in stellar populations across the galaxy and its companion. We found evidence for three distinct stellar populations, with ages of about 15~Gyr, 300~Myr and less than 10~Myr. We concluded that the interaction with the companion galaxy triggered star formation in both galaxies, and after a substantial time delay, the infalling gas triggered activity in the central supermassive black hole.  In order to better understand the connection between the radio source and its environment, we have undertaken further radio and optical observations of MRC~B1221$-$423, which we present in this paper. In sections~\\ref{sec:Radio-imaging} and~\\ref{sec:Radio-spectroscopy} we describe the radio images and \\hi\\ absorption spectra of MRC~B1221$-$423, and in section~\\ref{sec:Optical-observations} we describe the optical spectra. We have assumed a flat $\\Lambda$-CDM cosmology with $H_0 = 71\\;{\\rm km\\,s^{-1}\\,Mpc^{-1}}$, $\\Omega_\\mathrm{M}=0.27$ and $\\Omega_\\mathrm{\\Lambda}=0.73$\\ throughout. At the distance of the galaxy, 1~arcsec corresponds to a distance of 2.876~kpc.  \\begin{figure} \\centerline{\\psfig{figure=slitpos3.eps,width=\\columnwidth,clip=t}} \\caption{$V$-band NTT image of the host galaxy \\citep[from][]{jhcs05}, showing knotty features and tidal tails, strongly suggesting an ongoing interaction with the companion galaxy (arrowed). The 12~mm radio contours (Fig.~\\protect{\\ref{fig:contour_plot}}) are overlaid to show the scale of the radio source in relation to the galaxy. The two slit positions (Table~\\protect{\\ref{tab:obs}}) are shown; spectrum A was taken along the rotation axis of the galaxy (see \\S~\\ref{sec:Observ-data-reduct}), while spectrum B was taken along the line joining the galaxy to its companion, which is close to the axis of the radio source.} \\label{fig:slitpos} \\end{figure} ",
        "conclusions": "\\label{sec:discussion}  We have presented detailed radio and optical spectroscopic observations of MRC~B1221$-$423, a young CSS source which is currently undergoing a merger. Our main findings are \\begin{itemize} \\item a highly distorted radio morphology \\item a large Faraday depth and significant depolarisation in the northern, less distorted, lobe of the source \\item the presence of redshifted \\hi\\ in absorption \\item a range of stellar ages in the galaxy \\item strong emission lines with relatively low ionisation. \\end{itemize}  The radio observations of MRC~B1221$-$423 show dramatic interaction between the southern radio jet and the host galaxy, together with infalling gas. This jet appears to be bent through a full 180\\degr. The alignment of the magnetic field along the jet supports the interpretation of the feature as a single connected structure, with the magnetic field being oriented circumferentially where the extreme bending occurs. This is suggestive of a strong interaction with high density material. Interestingly, the RM synthesis shows no evidence of a dense interstellar medium where the jet bends.  The much shorter extent of the radio emission to the north, as well as the steep gradient in Faraday depth from north to south, suggests the presence of a dense and possibly highly magnetised medium local to the source. One possibility is that we could be seeing an orientation effect, with the northern jet pointed away from us and seen behind the galaxy. Large changes of jet direction in radio galaxies have been modeled as the result of a helical jet \\citep[see e.g.][]{hmcp84,cm93}; it is possible that a helix viewed almost head-on could explain the extreme change in projected jet direction that we see in MRC~B1221$-$423. The lack of evidence for a dense ISM at the bend of the jet from the RM synthesis would be consistent with this interpretation.  Alternately, the infalling gas could be interacting with the jet, producing both the change in jet direction and the different polarisation in the lobes.  \\citet{sg03} found a high degree of polarisation asymmetry in CSS sources, which they ascribed to interaction of the jets with infalling material.  \\hi\\ was clearly detected in absorption, with a narrow component at a velocity $+250\\kms$\\ compared to the velocity of the galaxy determined from the stellar component. Both this absorption and the radio continuum observations suggest the source is in a gas-rich environment. There is also a broader wing of \\hi\\ absorption extending to the systemic velocity.  The velocity shift with respect to the stellar redshift suggests large amounts of infalling gas. It should be noted, however, that the velocities of the \\hi, particularly of the broad component, are quite similar to the range of velocities shown by the emission lines (see Fig.~\\ref{fig:vprofile}), so it is possible the \\hi\\ absorption is just part of the overall distribution of gas. Redshifted \\hi\\ absorption is not common in similar sources: of 14 sources studied by \\citet{htm08}, most showed significantly \\textit{blueshifted} \\hi\\ absorption, reflecting outflow. Only two of their sources --- PKS~B0023$-$263 and PKS~B1934$-$638 --- showed redshifted \\hi\\ absorption.  The location of the absorbing gas is not clear. A comparison can be drawn with the radio galaxy NGC~315, which shows two absorption systems against the central region \\citep{mpo+09}.  In addition to the redshifted narrow component ($v \\sim +490$\\kms), NGC~315 has a broader (FWZI $\\sim 150$\\kms) component redshifted by $\\sim 80$\\kms\\ compared to the systemic velocity. \\citet{mpo+09} concluded from velocity gradients seen in VLBI data that the broad component arises from gas falling into the nucleus, while the narrow component arises from more distant gas clouds falling towards the galaxy. By analogy, we might be seeing a similar situation in MRC~B1221$-$423. The narrow absorption component could represent a discrete cloud at some distance from the galaxy, while the broad absorption arises much closer to the nucleus, either from a circumnuclear torus, or from the infalling gas that is fueling the nuclear activity. Higher resolution \\hi\\ observations are needed to distinguish these possibilities.  Optical spectra of MRC~B1221$-$423 reveal disturbed kinematics. Most of the ionised gas is at velocities similar to that of the stars, but in the nucleus we detect a broad blue-shifted component to the emission lines with $\\Delta v = 830$\\kms\\ FWHM, which we interpret as evidence for outflow of emission-line gas in the nucleus.  Since it is only visible in the nucleus, this could arise from gas interacting with the radio jet: the fact that it is blue-shifted might support the interpretation that the jet axis is strongly inclined to the line of sight.  The ionisation state of the gas is relatively low, with \\oii~3727 stronger than \\oiii~5007, and a very weak \\heii~4686 line. The line ratios change as a function of position, with the emission being dominated by the AGN close to the nucleus, and becoming more star-formation-like further out.  A similar result was found in the sample of compact radio sources studied by \\citet{htm09}, where the ionisation in the nuclear regions is dominated by the central AGN. Like them, we find a lower \\oiii~5007 luminosity than extended sources with similar radio luminosity in their comparison sample.  Star-formation activity is evident over the whole system, including the companion galaxy. The youngest stellar population is close to the nucleus, consistent with the idea that the same inflows of gas that have triggered star formation are also triggering the activity of the central black hole.  Our previous work on MRC~B1221$-$423 suggested the following scenario: tidal interaction with the companion galaxy triggered different phases of star formation in the two galaxies, and then, after a substantial delay, the central AGN was triggered to produce the radio emission. The current observations support this model. We have found evidence of past episodes of star formation, probable evidence for infalling gas, and indication of interaction between the radio jet and its environment.  Our work confirms the fact that young, powerful radio sources tend to live in a gas-rich nuclear environment. Further, we show that the AGN is younger than the starburst.  The ages of the stellar populations are comparable with the orbital period of the companion, which we estimate to be $\\sim 10^8$~yr, based on the projected separation of the companion galaxy from the nucleus, together with an estimate of the mass of the galaxy from the $K$-band luminosity. The faint arcs to the north-east \\citep{shp03} suggest there have been previous tidal interactions. These timescales suggest a time delay of $\\sim$~a few $\\times 100$~Myr between the tidal event (that triggered star formation) and the onset of AGN activity.  Such delays have been found in several nearby AGN \\citep[see e.g.][]{tds96,emt+06,dsm+07}, but this is possibly one of the best examples.  In recent work, \\citet{whc10} constructed a sample of starburst galaxies from the Sloan Digital Sky Survey. They found that the rate of accretion onto the central black hole, as measured by the strength of the \\oiii\\ line, rises steeply about 250~Myr after the onset of star formation. They speculate that the accretion onto the black hole may be being regulated by the starburst, with the accretion being dominated by low-velocity stellar winds, and accretion at early times being suppressed by supernovae.  We have also found indication of a possible infall of cold gas that could provide the fuel for the AGN. We note that the age distribution of the stellar populations is consistent with a picture in which cold or cooling gas is driven by the merger towards the centre of the host galaxy to fuel the AGN from relatively large radii, and that this gas begins to form stars while it is still well away from the nucleus. This contrasts with the situation in lower-luminosity radio sources in cluster environments, which, even though they sometimes contain the most massive black holes, are accreting in an environment dominated by hot gas and are associated with an accretion mode that produces low-power radio jets."
    },
    "1005/1005.2867_arXiv.txt": {
        "abstract": "{% A survey of mid-infrared gas-phase emission lines of \\htwo, \\water{} and various atoms toward a sample of \\nemb{} embedded low-mass young stars in nearby star-forming regions is presented. The sources are selected from the \\spitzer{} \"Cores to Disks\" (c2d) legacy program. }% {% The environment of embedded protostars is complex both in its physical structure (envelopes, outflows, jets, protostellar disks) and the physical processes (accretion, irradiation by UV and/or X-rays, excitation through slow and fast shocks) which take place. The mid-IR spectral range hosts a suite of diagnostic lines which can distinguish them. A key point is to spatially resolve the emission in the \\spitzer-IRS spectra to separate extended PDR and shock emission from compact source emission associated with the circumstellar disk and jets.}% {% An optimal extraction method is used to separate both spatially unresolved (compact, up to a few hundred AU) and spatially resolved (extended, thousand AU or more) emission from the IRS spectra. The results are compared with the c2d disk sample and literature PDR and shock models to address the physical nature of the sources.}% {% Both compact and extended emission features are observed. Warm (\\tex\\,few hundred\\,K) \\htwo, observed through the pure rotational \\htwo{} S(0), S(1) and S(2) lines, and \\si{} 25 $\\mu$m emission is observed primarily in the extended component. \\si{} is observed uniquely toward truly embedded sources and not toward disks. On the other hand hot (\\tex\\,$\\gtrsim700$\\,K) \\htwo, observed primarily through the S(4) line, and \\neii{} emission is seen mostly in the spatially unresolved component. \\feii{} and \\silii{} lines are observed in both spatial components. Hot \\water{} emission is found in the spatially unresolved component of some sources.}% {% The observed emission on $\\geq$1000 AU scales is characteristic of PDR emission and likely originates in the outflow cavities in the remnant envelope created by the stellar wind and jets from the embedded young stars. Weak shocks along the outflow wall can also contribute. The compact emission is likely of mixed origin, comprised of optically thick circumstellar disk and/or jet/outflow emission from the protostellar object.} ",
        "introduction": "\\label{sec:introduction}  Understanding the nature of deeply embedded protostars is a major challenge in the study of low-mass star formation. In its earliest stages, a young star is embedded in its natal envelope of cold gas and dust, prohibiting direct detection with optical and near-IR instruments. During the evolution of such protostars the envelope is dissipated partly due to the ongoing accretion onto the central star and disk and partly due to the action of the outflows driven by such stars. This embedded phase lasts only a short time, $\\sim6\\times10^5$ yr \\citep{evans09}, but is critical for the subsequent evolution of the star: it is the phase in which the final mass of the star, the overall mass of the circumstellar disk (and thus its ability to form planets), and the initial physical and chemical structure of the disk are determined.   The envelope surrounding young stars is subjected to highly energetic phenomena, in particular X-rays and UV radiation and streams of gas ejected from the star-disk boundary region. Both the slow and high velocity gas can shape the protostellar envelopes on a variety of spatial scales. The related physical processes, outlined below, have a strong impact on the physical structure and the molecular and atomic contents of the surrounding material. Mid-infrared spectroscopy, the central focus of this paper, provides a powerful tool to characterize and quantify the nature and energetics of these interactions as well as the composition and temperature structure of the material involved.  Based on previous infrared and submillimeter data, a general picture has emerged of the structure of embedded young stellar objects (YSOs) \\citep[e.g.,][]{young04,joergensen02,joergensen05,chiang08,tobin10}. Take as an example the isolated source HH\\,46\\,IRS (IRAS08242-5050) imaged with the Spitzer Space Telescope \\citep[see][and references cited]{noriega04,velusamy07} and studied in detail with submillimeter molecular lines as well as dust continuum emission \\citep[see summary in][]{vankempen09b}. The young star is surrounded by a roughly spherically symmetric envelope with a mass of $\\sim$5\\,M$_\\odot$ and a power-law density structure, with $n\\mathrm{(H_2)}\\,\\approx\\,10^6\\,\\mathrm{cm}^{-3}$ at 1000\\,AU from the star, decreasing to $\\sim 10^5$\\,cm$^{-3}$ at $\\sim$7000\\,AU. The infrared images reveal prominent red- and blue-shifted outflow lobes traversing the envelope with an opening angle up to 110$^o$ and continuing out to 0.2\\,pc. These outflow lobes are produced by jets and/or wide angle winds propagating at speeds up to 200\\,km\\,s$^{-1}$, which interact with the surroundings and create a shocked shell of material that is entrained and propelled outward at some speed \\citep[e.g.,][]{bally07,ray07,shang07}. The most prominent example of this interaction is the bow shock at the tip of the outflow. In addition, blobs of shocked warm gas and dust are seen at various locations along the outflow axis, which may be internal shocks in the (variable) wind or jet itself. Similar physical structures are expected for other embedded YSOs, although the densities at 1000\\,AU are up to an order of magnitude lower for sources with lower mass envelopes \\citep{joergensen02}.  The outflow cavities create a low-extinction pathway for ultraviolet (UV) photons produced in the accreting gas column at the star-disk boundary layer to escape out to much further distances \\citep{spaans95,bruderer09a,bruderer09b} than possible in a purely spherical symmetry \\citep[e.g.][]{stauber04}. For a UV luminosity of 0.1\\,L$_\\odot$ typical of low-mass YSOs, the enhancement of the UV radiation field with respect to the general radiation field is about \\hbox{$G_o$\\,$\\approx$\\,$10^4$} at 100\\,AU. Moreover, UV radiation is created along the outflow by the shocks themselves, especially by the higher velocity \\hbox{$J$-type} shocks \\citep{hollenbach89,neufeld89a}. Evidence for UV enhancements $G_o$ of a few hundred out to 0.1\\,pc have recently been inferred from observations of narrow high-$J$ CO submillimeter emission lines surrounding the outflow cavity walls \\citep[][see their Fig.\\,13]{vankempen09b,vankempen09c}. Thus, the Spitzer IRS slit of about $5-10''$ will pick up emission from gas with a range of densities from $10^4-10^7$\\,cm$^{-3}$, exposed to UV radiation with \\hbox{$G_o\\approx 10^2-10^4$}, and subjected to shocks with velocities up to several hundred km\\,s$^{-1}$ for typical cloud distances of \\hbox{$100-400$\\,pc}.   The consequences of UV radiation impacting on dense gas have been studied for decades through models of so-called Photon-Dominated Regions (or PhotoDissociation Regions, PDR) \\citep{tielens85a,tielens85b,hollenbach91,hollenbach99}. In the PDR layers closest to the UV source most molecular material is dissociated and the atomic species with ionization potential below 13.6\\,eV are ionized. This layer is characterized by bright fine structure lines such as \\oi, \\fei, \\feii{} and \\silii. \\si{} is never strong as most of the sulfur is in the form of S$^+$ in the warm PDR. At deeper layers, H$_2$ can survive, yet remains highly excited, resulting in strong pure rotational \\htwo{} lines \\citep[e.g.][applied to regions of massive star-formation]{kaufman06}.   Shocks impacting the surrounding envelope can be divided into two types: $J-$ and $C-$type shocks. The so-called dissociative \\jshocks{} \\citep{draine93,hollenbach97b,walmsley05} heat the gas up to $10^5$ K. Infrared emission comes from the dense hot post-shock gas in which most of the molecular material is still dissociated and a large fraction of the atoms are ionized resulting in strong emission lines of singly and doubly ionized atoms (e.g., \\feii{} 26 and 35 $\\mu$m, \\niii{} 6.6 $\\mu$m). This includes atoms which cannot be directly photoionized by photons with energies $<$13.6 eV such as the \\neii{} 12.8 $\\mu$m line \\citep{hollenbach89,hollenbach09}. Further downstream, molecular hydrogen forms in the cooling gas and strong high-$J$ pure rotational line emission is expected. \\neii{} emission is strong only in high velocity shocks \\hbox{($v_\\mathrm{s}>70$\\,km\\,s$^{-1}$)}, while the strength of other fine structure lines (e.g. \\si{} 25 $\\mu$m, \\fei{} 24 $\\mu$m and \\cli{} 11.4 $\\mu$m) does not strongly depend on the shock velocity \\citep[see][for details]{hollenbach89}. Note that these fast shocks imply travel times through the envelope of only ($<10^4$\\,yrs), unless the material is significantly slowed down by the interaction. For the so-called non-dissociative \\cshocks{} \\citep[e.g.][]{neufeld93,kaufman95,kaufman96}, the temperatures never become high enough to dissociate molecules; these shocks show strong lines from excited \\htwo, CO, \\water{} and OH \\citep[e.g.][]{kaufman96}.   To date, PDR and shock models have mostly been applied to the general interstellar medium (ISM) near massive young stars or supernova remnants. Previous observations of low mass young stellar objects (YSOs) using ground-based observatories or the ISO-SWS instrument \\citep[][]{degraauw96,vdancker_thesis} either lacked the sensitivity to observe the emission lines or the spatial resolution to separate the various emitting regions \\citep[see][for an overview]{vandishoeck04}. The sensitive InfraRed Spectrograph (IRS) \\citep[][]{houck04} on board the \\spitzer{} Space Telescope \\citep[][]{werner04} brings the detection of the mid-IR \\htwo{} pure rotational and atomic fine structure lines within reach for young Sun-like stars in nearby star-forming regions. The combination of high sensitivity, moderate spectral resolution $R=\\lambda/\\Delta\\lambda=600$, and modest spatial resolution makes \\spitzer{} well suited for the study of the gas in the ISM around low-mass young stars in the nearest ($\\lesssim300$\\,pc) clouds. Moreover, the \\spitzer{} aperture allows us, using an optimal extraction method, to disentangle both spatial components from the IRS echelle and long-slit spectral images down to scales of a few arcsec, better than can be achieved with single-dish millimeter telescopes. The combination of line fluxes, line ratios and spatial extent can help to identify the dominant heating processes -- PDR or shock-driven -- taking place, and thus the physical structure of the envelope on scales of a few hundred to a few thousand AU. The \\spitzer{} data thus complement traditional studies with (sub)millimeter telescopes.   Theory predicts that disks are established early in the embedded phase of star formation, as the inevitable byproduct of the collapse of a rotating core \\citep[e.g.][]{terebey84, cassen81}. Such 'young' disks, some of which show Keplerian velocity patterns, have been detected in submillimeter interferometry data \\citep[e.g.][]{keene90,lommen08,joergensen09}, but little is known about their physical and chemical structure. Detailed models of X-ray and UV irradiation of disk surface layers exist for more evolved pre-main sequence stars where the envelopes have been completely dissipated \\citep{bergin07,gorti08,woitke09,bergin09}. \\spitzer{} observations have revealed surprisingly rich spectra of such protoplanetary disks, with the lines seen either in absorption in a near edge-on geometry \\citep{lahuis06a} or in emission \\citep{carr08,salyk08}. Lines from \\water, HCN, \\acetylene, \\cotwo{} and OH have been detected, indicating hot molecular gas at temperatures of $\\sim500-1000$\\,K arising from the inner few AU of the disk. This hot gas is also seen in ground based CO infrared data \\citep[e.g.][]{najita03, blake04}. Such disks show various atomic transitions as well, most prominently the \\neii{} $12.8\\mu$m line \\citep{lahuis07c,pascucci06,espaillat07} \\citep[see also models by][]{meijerink08}. An interesting question is therefore whether young disks in the embedded phase studied here have similar spectra as the more evolved disks, or whether they show additional features, for example due to the accretion shock of material falling on to the disk \\citep{pontoppidan02,watson07}.   We present here an overview of the mid-IR gas-phase lines detected in embedded sources, including class 0 and class I sources observed in the \\spitzer{} legacy program ``From Molecular Cores to Planet Forming Disks'' (``Cores to Disks'' or c2d) \\citep[][]{evans03}. This program has collected a large sample of IRS spectra toward sources in the nearby Chamaeleon, Lupus, Perseus, Ophiuchus, and Serpens star-forming regions. High-$S/N$ spectra have been obtained within the $5-38\\,\\mu$m range for 226 sources at all phases of star and planet formation up to ages of $\\sim$5\\,Myr, of which 61 are presented here. Many of the sources are also included in the series of papers discussing the ice features in the c2d sample \\citep{boogert08,pontoppidan08,oberg08}. In Sections \\ref{sec:observations} and \\ref{sec:data-reduction} the source selection and data reduction are explained. In Section \\ref{sec:results} the observed atomic fine-structure and \\htwo\\ emission lines and the derived parameters are presented. Section \\ref{sec:discussion} reviews the results within the context of PDR and shock heating and by comparing the embedded sources in the present sample with the more evolved disk sources in the larger c2d program that have been discussed in \\citep[][]{lahuis07c}.  \\input f1.tex ",
        "conclusions": "A survey of the mid-infrared gas phase pure rotational lines of molecular hydrogen and a number of atomic fine structure transitions has been carried out toward a sample of \\nsources\\ class I sources with the \\spitzer-IRS of which \\nemb{} are truly embedded protostars. Both extended and compact source emission has been detected toward a majority of the sources. Compact emission, most prominently seen in hot \\htwo{}, \\neii{} and \\water{} lines, is seen toward $\\sim$1/4 to 1/3 of the truly embedded sources. The compact hot \\htwo{} and \\neii{} emission is believed to have its origin predominantly in dissociative \\jshocks{} in the outflow jets, though a true disk component may be present. The \\water{} emission is consistent with an origin in the inner regions of the young protostellar disk, similar to the \\water{} emission from the more evolved systems. Extended emission, dominated by low-$J$ \\htwo{} and \\feii{} emission, is observed toward $\\sim$2/3 of the truly embedded sources, both stage 0 and I sources. The extended emission is thought to originate from PDRs along the outflow cavity walls with possible contributions from non-dissociative \\cshocks.  These observations have shown the capability of \\spitzer{} to observe the gas around young stars. It clearly shows emission from PDR and shock excited gas by the presence of various \\htwo{} and atomic lines. However, its limited spatial and spectral resolution prevents a quantitative analysis of the excitation conditions.  Sources with atomic fine structure line emission, \\htwo{} and \\water{} emission are excellent candidates for follow up with various ground based instruments and near-IR and far-IR space-based instruments on \\herschel{} and JWST. The major far-IR cooling lines of [O\\,{\\sc i}] and [C\\,{\\sc ii}] as well as high-$J$ CO observable with \\herschel{} can provide important excitation constraints \\citep{vankempen10}. Several of the sources in this sample are targets in the ``Water in Star-forming regions with Herschel'' (WISH) and ``Dust, Ice and Gas in Time'' (DIGIT) \\herschel{} key programs. Future mid-IR observations with JWST-MIRI will allow the detection of weak emission lines and provide many more diagnostics to separate the physical processes. Such high spectral and spatial resolution observations are essential to constrain the location and extent of the emission (either by direct detection or through velocity resolved spectra) to properly understand the physical environment around the young stars.  The spectra of the compact component of the truly embedded sources show remarkable similarities with those of the evolved disk sources, both in terms of absolute fluxes and flux ratios. Both characteristic \\jshock{} emission, hot \\htwo{} and \\neii, and true disk emission through the molecular emission lines are observed. In particular, the molecular emission may offer a unique tool to study the young forming disks in their earliest phases and compare these with the more evolved stage II disks."
    },
    "1005/1005.0579_arXiv.txt": {
        "abstract": "The effects of astrophysical uncertainties on the exclusion limits at dark matter direct detection experiments are investigated for three scenarios: elastic, momentum dependent and inelastically scattering dark matter. We find that varying the dark matter galactic escape velocity and the Sun's circular velocity can lead to significant variations in the exclusion limits for light ($\\lesssim10$ GeV) elastic and inelastic scattering dark matter. We also calculate the limits using one hundred velocity distributions extracted from the Via Lactea II and GHALO N-body simulations and find that a Maxwell-Boltzmann distribution with the same astrophysical parameters generally sets less constraining limits. The elastic and momentum dependent limits remain robust for masses $\\gtrsim50$ GeV under variations of the astrophysical parameters and the form of the velocity distribution. ",
        "introduction": "Identifying the dark matter in our Universe is one of the great challenges facing both theorists and experimentalists. Dark matter direct detection experiments are now probing dark matter-nucleon cross sections expected of Weakly Interacting Massive Particles (WIMPs) and the results from the CDMS II \\cite{Ahmed:2009zw} and CoGeNT \\cite{Aalseth:2010vx} experiments may be hinting that a conclusive detection of dark matter is within reach of the next generation of experiments. The DAMA collaboration \\cite{Bernabei:2010mq} have also continued to measure an annual modulation signal, observed over the past thirteen years with an $8.9 \\sigma$ significance which has many of the properties expected from a dark matter origin. Although now inconsistent with the null results from other experiments assuming a vanilla elastically scattering WIMP, this persistent signal has inspired many models \\cite{TuckerSmith:2001hy, Bai:2009cd, Feldstein:2009tr, Chang:2009yt, Frandsen:2009mi, Masso:2009mu, Kopp:2009et,Essig:2010ye,Graham:2010ca} which can give rise to dramatically different experimental signals from the usual assumptions of elastically scattering spin-independent or spin-dependent interactions. Although these models may be capable of reconciling the DAMA result with other experiments, we consider them to be interesting, independent of the DAMA result, since they serve to highlight the different possible signals that could be produced at dark matter direct detection experiments.  As well as depending on particle physics models, the recoil spectrum and scattering rate depend on the astrophysical parameters and the form of the local velocity distribution the particle dark matter in the galactic halo takes. Unfortunately, these parameters still suffer from large uncertainties, while high resolution numerical N-body simulations have shown that the velocity distribution deviates significantly from the default choice, namely, a Maxwell-Boltzmann distribution.  Given this, in this paper we investigate the effect of astrophysical uncertainties on three scenarios: elastic spin-independent scattering; and, as proxies for the many varied models, inelastic and momentum dependent (aka form factor) spin-independent scattering. We vary the galactic escape velocity, the Sun's velocity relative to the galactic centre, and the local dark matter density within their experimental errors. We then compare the exclusion limits for these scenarios assuming a Maxwell-Boltzmann velocity distribution with those extracted from two of the highest resolution dark matter N-body simulations, Via Lactea II and GHALO. We particularly focus on four experiments with different target nuclei which span a large range of atomic number $A$ to compare the relative change at different experiments: CDMS II-Si; CDMS II-Ge; XENON10; and CRESST-II, with silicon ($A=28$), germanium ($A=72$), xenon ($A=131$) and tungsten ($A=184$) target nuclei respectively.\\footnote{Since we are seeking to understand the effects of the astrophysical uncertainties, we do not specifically address whether astrophysical uncertainties can improve the fit between the CoGeNT and DAMA signals. For analyses in this direction, see \\cite{Fitzpatrick:2010em,Chang:2010yk}.}  In Fig.\\ \\ref{3dRdE} we have plotted the differential scattering rate at a germanium detector to illustrate the differences between the three scenarios. Inelastic scattering of the dark matter to a state $\\sim100$ keV heavier \\cite{TuckerSmith:2001hy} changes the phenomenology in three principal ways: it leads to a suppression of scattering rates at low mass target nuclei experiments; an increase in the ratio of the modulated to unmodulated scattering rate; and a kinematical suppression of the recoil spectrum at low recoil energies. Although often invoked as a way of reconciling the modulation signal seen by DAMA with the null results at other experiments \\cite{TuckerSmith:2004jv, Chang:2008gd, MarchRussell:2008dy, Cline:2009xd, SchmidtHoberg:2009gn}, the latest (preliminary) results from CRESST-II \\cite{CRESSTWONDER} tentatively rule out this interpretation. However, having this splitting is natural in many models \\cite{Cui:2009xq,Alves:2009nf,Arina:2009um,ArkaniHamed:2000bq,MarchRussell:2004uf,Thomas:2007bu,MarchRussell:2009aq,ArkaniHamed:2008qn,Cholis:2009va,Gherghetta:2010cq} and leads to much other interesting phenomenology, for example through capture in astrophysical bodies \\cite{Nussinov:2009ft, Menon:2009qj, Shu:2010ta,McCullough:2010ai, Hooper:2010es}. Therefore, we emphasise that even if inelastic scattering can not reconcile the DAMA result with other experiments, it remains an interesting and viable possibility to explore.  In momentum dependent scattering \\cite{Feldstein:2009tr, Chang:2009yt}, the differential cross section is suppressed by additional powers of $q^2=2M_NE_R$, leading to a suppression of the exclusion limits from experiments which rely on having a sensitivity at low recoil energies, in our case XENON10 and CDMS II-Si. This dependence can arise for example through dipole or charge radius interactions \\cite{An:2010kc}.  In Section \\ref{Review} we briefly review the theory required to calculate scattering rates, then in Section \\ref{astrouncert} we review our current knowledge of the relevant astrophysical parameters and demonstrate and explain their effect on the experimental exclusion curves for elastic, inelastic and momentum dependent scattering. In Section \\ref{VLGHALO} we examine the effect of using velocity distributions from the Via Lactea II and GHALO simulations. Finally, Appendix \\ref{experiments} summarises the relevant experimental details, where in particular, we call attention to the reanalysed early CDMS II runs and discuss uncertainties affecting the XENON10 exclusion limits at low masses.  \\begin{figure}[t] \\centering \\includegraphics[height=2.0in]{1.pdf}\\caption{(Colour online). The differential scattering rate at a germanium detector for $M_{\\chi}=100$ GeV for elastic, inelastic and momentum dependent (MD) spin-independent scattering. The normalisation of each curve is different.} \\label{3dRdE} \\end{figure} ",
        "conclusions": "If we are to understand the nature of dark matter, it is vital that we have a full understanding of the astrophysical uncertainties affecting dark matter direct detection experiments. For spin-independent elastic and momentum dependent scattering with $M_{\\chi}\\gtrsim50$ GeV, we have shown that the exclusion limits are robust against variations in the galactic escape velocity $v_{esc}$ and the Sun's circular speed about the centre of the galaxy $v_0$ (right upper and left lower panels of Figs.\\ \\ref{v_esc} and \\ref{v0}) and under realistic variations in the form of the velocity distribution (middle and right panels of Fig.\\ \\ref{simelastic}), with uncertainty $\\sim 10\\%$. The major uncertainty in this mass range arises from the error in the local dark matter density (a factor of $2$).  In comparison, for lighter masses, we found that uncertainties in $v_0$ and $v_{esc}$ and the velocity distributions from the numerical simulations can shift the exclusion curves horizontally by $\\sim 1$ GeV at masses $M_{\\chi}\\sim10$ GeV (upper left panels of Figs.\\ \\ref{v_esc}, \\ref{v0} and \\ref{simelastic}). Similarly for inelastically scattering dark matter, we found the vertical shift in the exclusion curves when varying $v_0$ or $v_{esc}$ is large ($\\sim100\\%$), with lighter target experiments such as CDMS II being particularly affected. Variations in the velocity distribution can also lead to significant changes, as we explicitly demonstrated in Fig.\\ \\ref{siminelastic} for the CRESST-II experiment.  Inelastic and light dark matter are particularly sensitive to astrophysical uncertainties because the minimum speed needed to scatter is just below the galactic escape speed, so experiments only sample the tail of the velocity distribution. Even though the recoil spectrum for momentum dependent dark matter is different from elastically scattering dark matter (Fig.\\ \\ref{3dRdE}) they both respond to astrophysical uncertainties in a similar fashion for germanium, xenon and tungsten targets, because $v_{min}$ remains far from $v_{esc}$. Therefore models which only sample the tail of the velocity distribution should carefully examine the effect of astrophysical uncertainties on their limits."
    },
    "1005/1005.2783_arXiv.txt": {
        "abstract": "{We propose to regulate the infinities of eternal inflation by relating a late time cut-off in the bulk to a short distance cut-off on the future boundary.  The light-cone time of an event is defined in terms of the volume of its future light-cone on the boundary.    We seek an intrinsic definition of boundary volumes that makes no reference to bulk structures.    This requires taming the fractal geometry of the future boundary, and lifting the ambiguity of the conformal factor.  We propose to work in the conformal frame in which the boundary Ricci scalar is constant.  We explore this proposal in the FRW approximation for bubble universes. Remarkably, we find that the future boundary becomes a round three-sphere, with smooth metric on all scales. Our cut-off yields the same relative probabilities as a previous proposal that defined boundary volumes by projection into the bulk along timelike geodesics. Moreover, it is equivalent to an ensemble of causal patches defined without reference to bulk geodesics.  It thus yields a holographically motivated and phenomenologically successful measure for eternal inflation.}  \\begin{document} ",
        "introduction": "String theory gives rise to an enormous multiverse where the constants of nature are locally constant but vary over extremely large distance scales~\\cite{BP,KKLT}. In this context, it is natural to make predictions by counting. The relative probability of two events (for example, two different outcomes of a cosmological or laboratory measurement) is defined by their relative abundance, \\begin{equation} {p_1 \\over p_2} = {N_1 \\over N_2}~, \\end{equation} where $N_1$ is the expectation value of the number of times an event of type 1 occurs in the multiverse.  There is ambiguity in computing the ratio $N_1/N_2$: naively, both numbers are infinite.  Starting from finite initial conditions, eternal inflation produces an infinite spacetime volume in which everything that can happen does happen an infinite number of times.  A procedure for regulating this divergence is called a measure.  Different measures can lead to different relative probabilities starting from an otherwise identical theory. \\begin{figure} \\begin{center} \\includegraphics[scale = .7]{lightconecut} \\end{center} \\caption{Constant light-cone size on the boundary defines a hypersurface of constant ``light-cone time'' in the bulk.  The green horizontal lines show two examples of such hypersurfaces. They constitute a preferred time foliation of the multiverse. In the multiverse, there are infinitely many events of both type 1 and type 2 (say, two different values of the cosmological constant measured by observers). Their relative probability is defined by computing the ratio of the number of occurrences of each event prior to the light-cone time $t$, in the limit as $t\\to \\infty$.} \\label{fig-lightconecut} \\end{figure}  A number of measures have been proposed, including~\\cite{LinMez93,LinLin94, GarLin94,GarLin94a,GarLin95, GarSch05,VanVil06,Van06,Bou06,Lin06,Lin07,Pag08, GarVil08,Win08a,Win08b,Win08c,LinVan08,Bou09}, but over the past several years a vigorous phenomenological effort (e.g., \\cite{GarVil05,FelHal05,PogVil06,FelHal06,SchVil06,Sch06,BouHar07,BouYan07,OluSch07,BouFre07,Sch08,BouLei08,Pag06b,HalNom07,HalSal07a,HalSal07b,DGSV08,BouFre08b,DGLNSV08,HalSal08,Sal09,BouHal09,PhiAlb09,BouLei09,HalNom09,EloGoh09}) has focused attention on a few simple proposals that remain viable. In this paper, we will focus on the light-cone time cut-off, which arises from an analogy with the AdS/CFT correspondence.  Given a short-distance cut-off on the boundary conformal field theory, the radial position of the corresponding bulk cut-off~\\cite{SusWit98} can be obtained from causality alone~\\cite{BouRan01}, without reference to the details of the bulk-boundary correspondence.  In the multiverse, the time of a bulk event can similarly be defined in terms of a scale on the future boundary of the multiverse~\\cite{GarVil08}.  The simplest causal relation is to associate to each bulk event a boundary scale given by the volume of its future light-cone on the boundary~\\cite{Bou09} (see Fig.~\\ref{fig-lightconecut}).\\footnote{A different bulk-boundary relation was proposed in Ref.~\\cite{GarVil08}.  One of us (RB) has argued that this relation is less well-defined than light-cone time and is not analogous to the UV/IR relation of AdS/CFT~\\cite{Bou09}.  These concerns aside, it could be combined with the metric we construct on the future boundary.  The resulting measure would be different from the one obtained here.}  There is a remaining ambiguity: how is the volume of the future light-cone to be defined?  Near the future boundary, physical distances diverge in inflating regions and go to zero near singularities. Ref.~\\cite{Bou09} defined boundary volumes by erecting a congruence of geodesics orthogonal to a fixed, fiducial bulk hypersurface and projecting future infinity onto the fiducial hypersurface along the geodesics. The definition of Ref.~\\cite{Bou09} is quite robust\\footnote{In this respect, the definition of volumes explored here remains inferior, for now, since it is completely well-defined only for homogeneous bubble universes.  However, for the reasons stated below, it may ultimately prove to be more fundamental.}; in particular, it does not matter whether geodesics expand or collapse or cross.  We would like to find an alternative definition of the boundary volume which does not rely on such an elaborate bulk construction. The geodesics used to project the light-cone onto the fiducial hypersurface do only one thing for us: they define boundary volumes. Yet they encode an enormous amount of geometric bulk information (the exact path of each geodesic, the expansion and shear of nearby geodesics, etc.), most of which is never used for any purpose.  They bear no apparent relation to any physical system; for example, they do not represent the worldlines of actual particles.  Moreover, the construction takes an absurdly classical viewpoint of the bulk geometry: because of the exponential expansion of de~Sitter vacua, the projection of a late time light-cone onto the fiducial hypersurface has subplanckian volume.  There is another reason why it would be nice to eliminate bulk geodesics from the definition of the measure.  Ultimately, one expects that a fundamental description of the multiverse will involve its boundary structure in some way. This was a key motivation for seeking a multiverse analogue of the UV/IR relation of AdS/CFT in the first place.  The equivalence of the causal patch and light-cone time cut-off, along with their phenomenological successes, encourages us to take seriously the motivations behind the two measures.  In particular, we would like to define the light-cone time cut-off, to the greatest extent possible, in terms of quantities that are intrinsic to the boundary.  In this paper we will propose an intrinsic definition of volumes on the future boundary of the multiverse, and we will explore the resulting time foliation and measure.  Like in AdS, we will render the boundary finite by a conformal rescaling of the bulk metric.  However, we face two difficulties that have no direct analogue in AdS/CFT. \\paragraph{The boundary of the multiverse is naturally a fractal.}  In the bulk, bubbles of different vacua keep being produced at later and later times, leading to boundary features on arbitrarily small scales.  If the bulk evolution exhibits attractor behavior (this is common to many global foliations of the multiverse, though the attractor itself depends on the foliation), then the boundary will exhibit self-similarity in the short-distance limit.  Because of the abundance of short-distance structure, the boundary must be constructed first on large scales, and then refined to better and better resolution, corresponding to evolving the bulk to later and later times.  The details of this procedure have not been carefully formulated and will concern us greatly.  The fractal structure leads to a number of difficulties for defining a metric on the future boundary: \\begin{figure}\\label{fig-onebubblemultiverse} \\begin{center} \\includegraphics[scale = .35]{OneBubbleMultiverse} \\end{center} \\caption{Conformal diagram of de~Sitter space with a single bubble nucleation.  The parent de~Sitter space is separated from the daughter universe by a domain wall, here approximated as a light-cone with zero initial radius (dashed line).  There is a kink in the diagram where the domain wall meets future infinity.  This diagram represents a portion of the Einstein static universe, but the Ricci scalar of the boundary metric is not constant.} \\end{figure} \\begin{itemize}  \\item Generically, when a new bubble forms at late times in the bulk, one would like to include this information on the boundary, for example by coloring a small disk on the boundary which is enclosed by this bubble.  But the new vacuum will also change the metric in the bulk, and thus, the shape of the conformal diagram.  In other words, generically, the boundary {\\em metric\\/} will keep changing as we increase the resolution.  This is worse than the behavior of most fractals we are familiar with, which are defined by coloring points on a {\\em fixed} background metric.  \\item A natural way of constructing a conformal diagram depicting the formation of a bubble universe in de~Sitter space leads to a future boundary with a ``sharp edge'' (see Fig.~2), where the first derivative of the induced boundary metric is discontinuous.  With more and more bubbles forming at late times, such edges would appear on arbitrarily small scales.  \\end{itemize}     \\paragraph{The shape of the boundary is defined only up to finite conformal transformations.}  Because the boundary metric is obtained from the bulk spacetime by a conformal rescaling, it is only defined up to a finite conformal factor. The shape of many conformal diagrams can be changed by a conformal transformation.  The resulting diagram is just as legitimate as the original one, but its boundaries may have a different geometry.  \\begin{itemize} \\item In the multiverse, this ambiguity in the boundary metric leads to an ambiguity in the bulk foliation.  Unlike in AdS, which is asymptotically empty, this leads to a potential ambiguity in the measure.  For example, a different choice of conformal factor may change the predicted probability distribution for the cosmological constant. \\end{itemize}   \\paragraph{Outline} We introduce the concept of light-cone time in Sec.~\\ref{sec-lct}, emphasizing that it requires a definition of the volume of any future light-cone.  In Sec.~\\ref{sec-ricci}, we propose that the volume should be defined as the volume enclosed by the light-cone on the future conformal boundary of the multiverse.  Moreover, we propose choosing the conformal factor so that the boundary metric has constant scalar curvature, \\begin{equation} R = {\\rm constant}~, \\label{eq-rcon} \\end{equation} where $R$ is the Ricci scalar of the boundary metric\\footnote{Garriga and Vilenkin \\cite{GarVil09} defined the metric on future infinity by foliating the spacetime  by surfaces whose induced metric has $R = 0$. In general spacetimes this is different from our proposal, and bulk $R=0$ surfaces do not always exist. However, in the special case we focus on of homogeneous de Sitter bubbles, their boundary metric is the same as ours.}  A recent result in mathematics---the solution to the Yamabe problem---guarantees that a suitable conformal transformation can always be found and is essentially unique on the boundary.  This proposal defines what we shall refer to as ``new light-cone time''.  In Sec.~\\ref{sec-frw}, we apply our proposal to a simplified landscape model.  We approximate bubble universes as homogeneous open FRW cosmologies, we assume that they all have nonzero cosmological constant, and we neglect bubble collisions.  Despite these simplifications, all of the difficulties listed above arise in this model, and we show that our proposal succeeds in addressing them.  We construct the conformal diagram iteratively, making sure that the boundary condition $R=$ const.\\ is satisfied each time a new bubble universe is included in the bulk. We find that this procedure leads to a fixed, everywhere smooth boundary metric that can be taken to be a unit three-sphere.  The boundary metric picked out by the condition $R=$ const.\\ defines new light-cone time. In Sec.~\\ref{sec-prop}, we present some of the most important properties of this time foliation.  In Sec.~\\ref{sec-rate}, we derive rate equations that govern the distribution of vacua.    We emphasize the boundary viewpoint, in which bubble universes correspond to topological disks on the boundary three-sphere. In Sec.~\\ref{sec-asol}, we find the solution of the rate equations.  The late time attractor behavior in the bulk corresponds to a universal ultraviolet scaling of the distribution of disks on the boundary.  In Sec.~\\ref{sec-count}, we derive the crucial expression that underlies the probability measure: the number of events of arbitrary type, as a function of light-cone time.  In Sec.~\\ref{sec-pm}, we analyze the probability measure.  In Sec.~\\ref{sec-gloglo}, we show that new light-cone time yields the same measure as old light-cone time, i.e., both cut-offs predict the same relative probabilities for any two types of events in our simplified landscape model.  In particular, this implies that our ``new'' light-cone time shares the phenomenological successes of the old one, and of the causal patch measure dual to it.  In Sec.~\\ref{sec-general}, we discuss how our approach may extend to the general case, where inhomogeneities, collisions, and vacua with $\\Lambda=0$ are included.  We also consider the (likely) possibility that the landscape contains vacua of different dimensionality.  In defining a unique boundary metric, several difficulties arise in addition to the ones listed above, and a more general method of implementing Eq.~(\\ref{eq-rcon}) is needed.  We discuss what phenomenological properties may be expected of the resulting measure.  In particular, we expect that in the context of inhomogeneous universes, a boundary definition of light-cone volume will address a problem pointed out by Phillips and Albrecht~\\cite{PhiAlb09}. ",
        "conclusions": ""
    },
    "1005/1005.0924_arXiv.txt": {
        "abstract": "We give an estimation of the masses of light dark matter particle and dark energy quasiparticle which can be extracted from the astrophysical data about the contributions of baryon ($\\Omega_{B}$), dark matter ($\\Omega_{DM}$), and dark energy ($\\Omega_{X}$) densities to the total matter-energy density budget in our universe. We use the quantum cosmological model in which dark energy is a condensate of quantum oscillations of primordial scalar field. The dark energy quasiparticle with the mass $\\approx 15$ GeV is consistent with  the 7-year WMAP and other data on $\\Omega_{B}$ and $\\Omega_{X}$. The quasiparticles can decay with violation of CP-invariance into baryons, leptons, and dark matter. The WIMP mass $\\approx 5$ GeV corresponds to observed values of $\\Omega_{B}$ and $\\Omega_{DM}$. Such a mass agrees with the observations of CoGeNT, DAMA, and CDMS. Quasiparticles of dark energy can be registered as a constant background of radiation with the frequency $\\approx 3.64 \\times 10^{24}$ s$^{-1}$. Dark matter particles must exhibit themselves in the form of signals with the frequency  $\\approx 1.21 \\times 10^{24}$ s$^{-1}$ of  radiation from galaxy clusters where the intensive gravitational fields produced by dark matter exist. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.5730_arXiv.txt": {
        "abstract": "Using effective field theory (EFT) techniques we calculate the next-to-leading order (NLO) spin-orbit contributions to the gravitational potential of inspiralling compact binaries. We use the covariant spin supplementarity condition (SSC), and explicitly prove the equivalence with previous results by Faye et al. in \\url{gr-qc/0605139}. We also show that the direct application of the Newton-Wigner SSC at the level of the action leads to the correct dynamics using a canonical (Dirac) algebra. This paper then completes the calculation of the necessary spin dynamics within the EFT formalism that will be used in a separate paper to compute the spin contributions to the energy flux and phase evolution to NLO. ",
        "introduction": "LIGO \\cite{ligo}, VIRGO \\cite{virgo} and eventually LISA \\cite{lisa} and Einstein Telescope \\cite{etele} expect to detect radiation from inspiralling binary systems, and building templates for these events has become increasingly important \\cite{gwaves1,gwaves2}. While for late stages of the inspirals numerical techniques are needed \\cite{luis,frans,carlos}, for the early stages we may rely on the Post-Newtonian (PN) approximation as an expansion in small velocities $(v/c)\\ll 1$ (see \\cite{review} for a review and further references.)  As argued in the literature  \\cite{pan, buo3, hughes,arun}, the expectation that black holes may be close to maximally rotating in the binaries \\cite{McClintock} produces a significant impact in the accuracy of gravitational wave templates. Therefore, if not for detection, parameter extraction justifies the need for physical templates that include spin contributions at higher PN orders beyond the leading effects. In addition, PN corrections including spin are relevant for comparison between analytic results and numerical simulations \\cite{lousto}.\\\\  In order to produce such templates an important building block is the gravitational potential, $V$, responsible for the dynamics of the bodies in the binary system. The leading order (LO) spin-orbit and spin-spin contributions to $V$ have been known in the literature for quite some time \\cite{wald, barker,dass, ehlers, kidder1, kidder2,poisson}, and spin precession has already been observed in some binary pulsars, e.g. \\cite{prec}. On the other hand, only recently the next-to-leading order (NLO) spin-orbit \\cite{owen, bbf, damour}, spin(1)spin(2) \\cite{eih,proc, comment,nrgrss, Schafer3pn} and spin(1)spin(1) \\cite{nrgrs2,Schafer3pn2,Hergt} contributions have been computed, albeit with rather dissimilar techniques.  The NLO spin-orbit effects (2.5PN) were computed in \\cite{owen, bbf} within the `standard' PN approach \\cite{review}, namely by obtaining the metric using Bailey-Israel-Dixon's spin-dependent stress energy tensor \\cite{dixon} and then solving Mathisson-Papapetrou equations \\cite{papa}. On the other hand, in \\cite{damour} the frequency of spin precession was directly obtained in ADM coordinates by a `suitable redefined constant-magnitude spin vector' and subsequently a Hamiltonian was derived that was shown to be equivalent to the results in \\cite{bbf}. This is a very useful trick that we borrow in this paper, although we will obtain the frequency of precession using a completely different approach.\\\\  The NLO spin(1)spin(2) effects (3PN) were first obtained in \\cite{eih,proc,comment} using a new formalism based on the application of Effective Field Theory (EFT) techniques introduced in \\cite{nrgr}, where it was coined Non-Relativistic General Relativity (NRGR), and extended to include spinning constituents in \\cite{nrgrs}. An EFT framework\\footnote{See \\cite{iraeft} for a review on EFT techniques.} has shown to be extremely powerful in many realms of {\\it classical} gravitational (and non-gravitational) physics. For example, it has been used to compute the NNLO spin-independent contribution to the potential (2PN) \\cite{andi1}; absorption effects for compact objects in binary systems and time dependent backgrounds \\cite{dis1,dis2}; the electromagnetic \\cite{adam} and gravitational self-force on extended objects \\cite{chad1}; radiation-reaction effects in the extreme mass ratio limit \\cite{chad2}; corrections to thermodynamic quantities for caged black holes \\cite{cbh,andi2}; deviations from General Relativity at the nonlinear level \\cite{mag}; and recently applied to the study of cosmological perturbations \\cite{eftfluid}.  The radiation sector of NRGR has been developed further in \\cite{rad1, spinrad}. The purpose of this paper is thus to compute the NLO spin-orbit corrections to the gravitational potential\\footnote{See \\cite{delphine} for an alternative derivation of the NLO spin-orbit potential.} (within our set of conventions), which together with the spin-spin results in \\cite{eih,comment,proc,nrgrss,nrgrs2} and source multipole moments computed in \\cite{rad1, spinrad} allow us to obtain all the spin contributions in the energy flux and phase evolution to 3PN order. This will be reported in a separate paper\\footnote{A comparison with the recent results for the NLO spin-orbit contributions to the radiated power in \\cite{buo2} will be performed in \\cite{spinrad2}.} \\cite{spinrad2}.\\\\  A crucial aspect of the calculations dealing with spin in General Relativity is the choice of spin supplementarity conditions (SSCs) \\cite{pryce,tulc} and, in a Lagrangian/Hamiltonian framework, the resulting (Dirac) algebra in the reduced phase space. This algebra is expected to be quite cumbersome in generic SSCs, already in a Minkowski background \\cite{regge}. Nevertheless, the results in \\cite{eih,comment} were obtained using the Newton-Wigner (NW) SSC at the level of the action, which was shown to be a correct procedure in \\cite{proc} using standard EFT power-counting techniques. As we shall show in this paper, the NW SSC also leads to a canonical structure at linear order in spin, at least to NLO. This is consistent with the results in \\cite{brb} (although formally the comparison only applies in the test-particle limit.)  Subsequently the potential in \\cite{eih,comment} was re-derived using a Routhian approach, introduced in \\cite{proc}, where the (covariant) SSC is only imposed at the late stages of the calculation so that we may never depart from a canonical algebra, and full consistency was found \\cite{nrgrss}. The NLO spin(1)spin(2) Hamiltonian has also been computed in \\cite{Schafer3pn} in ADM coordinates, and the equivalence with the results in \\cite{eih} was shown in \\cite{comment}.\\\\  The basic idea of our approach is the systematic separation of the relevant scales $(\\lambda_{\\rm fs} \\ll \\lambda_{\\rm p} \\ll \\lambda_{\\rm r})$ in the two body problem; from finite size effects ($\\lambda_{\\rm fs} \\sim Gm$ for compact objects), to the potential ($\\lambda_{\\rm p} \\sim r$) and gravitational radiation ($\\lambda_{\\rm r} \\sim \\frac{r}{v}$) scales (for a review see \\cite{nrgrLH}.) For the theory of potentials the NRGR prescription consists on computing all possible (Feynman) diagrams (without external radiation gravitons) scaling with a definite power of $v$. For spinning extended objects, the `rules' are derived from an action (Routhian) principle and are listed in \\cite{nrgr,nrgrs,nrgrss,nrgrs2}. The EFT formalism is therefore significantly more efficient than the (more traditional) PN methods since it maps complex integrals into the computation of Feynman diagrams, uses textbook regularization techniques, and it is especially suited to handle spin degrees of freedom. Feynman diagrams allow for a very natural systematization and (physical) visualization of the computation, and moreover it can be automatized almost entirely using {\\it Mathematica} code \\cite{nrgrwork}. As a result, the calculation of the NLO spin(1)spin(2) potential was no more involved than obtaining the Einstein-Infeld-Hoffmann Lagrangian as in \\cite{nrgr}, contrary to the intricate approach of \\cite{Schafer3pn}.\\\\  The NLO spin(1)spin(1) potential (also at 3PN) was computed in \\cite{nrgrs2}. Most of the spin(1)spin(1) effects are due to (self-induced) finite size contributions \\cite{nrgrs} (and an extra term responsible for the preservation of the SSC \\cite{proc,nrgrs2}.) In the spirit of EFTs these terms are encoded in higher dimensional operators in the effective action (constrained only by the symmetries of the theory) up to unknown coefficients that are `matched' with a given observable in the complete theory \\cite{nrgr}, for example the quadrupole moment of a spinning black hole \\cite{nrgrs,nrgrs2}. Once again the gravitational potential follows as a sum of Feynman diagrams of definite PN order which greatly simplifies the treatment of finite size corrections. These results were corroborated in \\cite{Schafer3pn2, Hergt}.\\\\  The Hamiltonian of \\cite{Schafer3pn2,Hergt} was partially obtained by imposing Poincare invariance, which is not sufficient to determine all the coefficients, and the extra terms were fixed `by a three-dimensional covariant ansatz for the source terms of the Hamilton constraint.' As argued in \\cite{Schafer3pn2}: `the arguments that lead to a unique fixation of the coefficients are, however, far from being straightforward.'  On the other hand, the determination of the potential in \\cite{nrgrs2} is straightforward once a (very natural) higher dimensional operator is included. Moreover, using EFT power counting techniques one can easily show that one (and only one) operator is sufficient to NLO \\cite{nrgrs, nrgrs2}, whose overall (Wilson) coefficient is directly related to the spin quadrupole of the Kerr black hole.  Truth be told, the advantage of the Hamiltonian in \\cite{Schafer3pn2} is that the spin tensor is fully reduced, and one only deals with a three-vector with canonical brackets. In our formalism we bypass the complications of the (Dirac) algebra in the covariant SSC by reducing to a three-vector as the final step in the calculation. Therefore, in the intermediate steps we have an extra three-vector degree of freedom, ${\\bf S}^{(0)i} \\equiv S^{0i}$, with non-vanishing brackets (see Eq. (\\ref{als0})), that is only solved for via the SSC once the equations of motion are obtained (see Eqs. (\\ref{alssc} - \\ref{spind}).) This however does not represent any major difficulty, no more than dealing with an antisymmetric spin tensor obeying a canonical algebra\\footnote{One can also perform a transformation that reduces the spin 3-vector algebra to standard brackets \\cite{Hergt}.}.\\\\  Ultimately, given the differences in approaches, the agreement found in all cases gives us strong confidence the results are correct. In this paper we add yet another piece of evidence for the validity of the spin potentials by computing the NLO spin-orbit effects using the EFT framework of NRGR. The computation proceeds systematically as a sum of Feynman diagrams. Hence first we sketch\\footnote{We encourage the reader to consult the NRGR literature for more details \\cite{nrgr,nrgrs,eih,proc,nrgrLH,nrgrss,nrgrs2}.} the ingredients that determine the necessary Feynman rules before we embark in the computation of the diagrams, that we divide into terms linear and quadratic in $G$ (Newton's constant.) Divergences are handled by textbook regularization techniques that result in finite integrals. After the potential is obtained we perform a spin redefinition that takes the spin dynamics into precession form from which we  read off the precession frequency. The details are summarized in an appendix. Then it is straightforward to show that our result resembles those in \\cite{bbf} since both frequencies, ours and the one given in \\cite{bbf}, are related by a total time derivative. Following \\cite{damour} one can then construct a Hamiltonian whose equations of motion turn out to be canonically related to those in \\cite{bbf}, and the equivalence is thus proven.  In the last section we show that imposing the NW SSC at the level of the action also leads to the correct dynamics in a canonical manner. Remarkably, the transformation found in \\cite{regge} between covariant and NW SSCs in a flat background applies to our case once re-written in a locally flat frame and with derivatives taken with respect to the `free falling' proper time. Therefore, here we extend the results in \\cite{brb} to the case of self-gravitating objects, as long as we ignore spin(1)spin(1) (finite size) terms\\footnote{Imposing the NW SSC does not lead to canonical brackets in the spin(1)spin(1) sector since it knocks off a necessary extra term in the Routhian, proportional to the SSC, which guarantees its preservation upon time evolution \\cite{proc,nrgrss,nrgrs2}.}. ",
        "conclusions": "Using the EFT formalism of NRGR \\cite{nrgr} for spinning extended objects \\cite{nrgrs,proc}, we computed the NLO contributions to the spin-orbit gravitational potential of inspiralling compact binaries. We explicitly proved the equivalence with previous results in the literature \\cite{bbf,damour} and provided extra support for the canonical aspects of the NW SSC (that was already shown to be the case for the spin(1)spin(2) sector to 3PN order in \\cite{proc}.) This is consistent with the claims in \\cite{brb} and here we extended it to the case of self-gravitating objects, as long as we ignore finite size effects. Moreover, we showed that we can adapt the transformation between NW and covariant SSCs found in \\cite{regge} for a spinning particle in a flat background to the (gravitationally bound) binary system.  In a separate paper the results reported here and in \\cite{eih,comment,nrgrss,nrgrs2,rad1,spinrad} will be used to compute the spin-orbit and spin-spin contributions to the energy flux and phase evolution to 3PN order \\cite{spinrad2}.  \\begin{center} {\\bf Acknowledgments} \\end{center}  I would like to thank Delphine Perrodin and Andreas Ross for discussions and checks on Feynman diagrams and potential. This work was supported in part by NSF under Grant No.04-56556.  \\appendix"
    },
    "1005/1005.3116_arXiv.txt": {
        "abstract": "\\noindent We perform a detailed and quasi model-independent analysis of direct annihilation of Dark Matter into neutrinos. Considering different cases for scalar and fermionic Dark Matter, we identify several settings in which this annihilation is enhanced, contrary to some statements in the literature. They key point is that several restrictions of, e.g., a supersymmetric framework do not hold in general. The mass generation mechanism of the neutrinos plays an important role, too. We illustrate our considerations by two examples that are not (as usually) suppressed by the smallness of the neutrino mass, for which we also present a numerical analysis. Our results can be easily used as guidelines for model building. ",
        "introduction": "\\label{sec:intro}  Astrophysics and cosmology provides us with compelling evidences of a form of non-luminous and non-baryonic matter, the so-called Dark Matter (DM), which should account for almost 23\\% of the total energy density of the Universe. One possibility to detect DM is to search for neutrinos coming from annihilations of DM particles in the Milky Way galactic center and in the galactic halo, in dwarf spheroidal galaxies, and in celestial bodies, like the Earth and the Sun. For recent studies on this subject see, e.g., Refs.~\\cite{Hooper:2008cf,Feng:2008qn,Kumar:2009ws,Niro:2009mw}. The main scope of this paper is to analyze the DM pair-annihilation into neutrino final states as model-independent as possible. The most interesting signal to look for at neutrino telescopes is represented by a monochromatic neutrino signal, which can be produced by DM pair-annihilations directly into $\\nu \\bar{\\nu}$ (or into $\\nu \\nu$ ($\\bar{\\nu}\\bar{\\nu}$), if we allow for lepton number violating (LNV) processes). The neutrino energy spectra produced in these annihilation channels are constituted of a soft part (originating from higher order processes) and a sharp line at an energy $E_\\nu \\simeq m_\\chi$. We restrict our study to the two-body direct production, since this is the golden channel for DM discovery at neutrino telescopes. In our analysis, we distinguish between Dirac and Majorana neutrinos and, in the latter case, we consider different neutrino mass generation mechanisms, since they affect the relation between the physical neutrino mass and the neutrino Yukawa couplings.  In the literature, it is often stated that direct annihilation into neutrinos is always suppressed. Though this statement is perfectly true in supersymmetric models (see, e.g., Refs.~\\cite{Hooper:2008cf,Goldberg:1983nd,Jungman:1995df}), it does not hold in general. The key point of our paper is to investigate the cases, where such a suppression is not present. This can already be seen in a simple see-saw framework (using only one generation of $\\nu_L$ and $N_R$): Here, it is often correctly argued that, even though the neutrino Dirac Yukawa coupling $y_D$ might be sizable, it still does not allow for a sizable direct production of two (light) neutrinos, since it couples only to one doublet $\\nu_L$ as well as to one singlet interaction eigenstate $N_R$. Even though the doublet state practically consists of the light mass eigenstate only, $\\nu_L\\approx \\nu'_L$, the fraction of light mass eigenstates within the singlet $N_R$ is suppressed by a tiny mixing angle $\\theta\\sim \\frac{m_D}{M_R}$, which reduces the corresponding production rate. However, when going to a see-saw type II framework, the situation changes: For illustration, one could add a mass term $m_L \\overline{(\\nu_L)^C} \\nu_L$, whose mass is given by the scalar triplet vacuum expectation value (VEV) $v_T$ times the corresponding Yukawa coupling $y_T$. The VEV $v_T$ is forced to be small, due to the correction to the $\\rho$-paremeter, which allows $y_T$ to be larger. But $y_T$, in turn, will be related to the annihilation rate of DM into neutrinos in an $s$-channel diagram, in case the DM can couple to the Higgs triplet. Furthermore, the left-handed mass $m_L$ is related to the physical neutrino mass by $m_\\nu=m_L-\\frac{m_D^2}{M_R}$ (see Sec.~\\ref{sec:Neutrino:mass:see-saw} for details), which implies that a partial cancellation in the see-saw formula might be responsible for the correct neutrino mass and allows for an $m_L$ that could be much larger than $m_\\nu$. Another example is if the triplet does not obtain a VEV at all: Then, there will be no restriction on the size of $y_T$ from the neutrino sector and the annihilation rates could be even larger. Note that, however, in this case other bounds may apply, e.g.\\ from limits on lepton flavour violation~\\cite{Kakizaki:2003jk} or neutrino-less double beta decay~\\cite{Petcov:2009zr}. This shows how the suppression of the annihilation rates into neutrinos is not always realized, and we will analyze all possibilities in a model-independent way.  This paper is structured as follows: After a brief review of the possible neutrino mass terms in Sec.\\,\\ref{sec:Neutrino:mass}, we report in Sec.\\,\\ref{sec:nuprod} the various possibilities for monoenergetic neutrino production, considering explicitly scalar and fermionic DM, as well as the corresponding $s,t,$ and $u$ channels. For simplicity, we do not extend the Standard Model (SM) gauge group, but we contemplate different $SU(2)_L$ representations for the DM and the mediator particles. This kind of systematic analysis was not presented before in the literature in the context of DM annihilation into neutrinos (see Ref.~\\cite{Agrawal:2010fh} for the classification of DM scattering off nucleons dominated by spin-dependent interactions). Our results are summarized in Sec.\\,\\ref{sec:general}, before we give a discussion of examples of unsuppressed scenarios in Sec.\\,\\ref{sec:Unsuppressed}. We explicitly show the behavior of the annihilation cross sections for a promising $s$-channel and $t$-channel diagram, considering both, the case of scalar and fermionic DM. Our results are then compared to experimental limits on $\\mu$ and $\\tau$ decays, as well as on lepton flavour violating (LFV) processes. The constraints coming from neutrino searches are also considered. All cross sections we have used are listed in Appendix~A. For specific models in which the DM particles annihilate mainly in neutrinos see, e.g., Refs.\\,\\cite{Falkowski:2009yz,Blennow:2009ag}. ",
        "conclusions": "\\label{sec:conc}  We have performed a detailed model-independent analysis of the process of Dark Matter annihilation directly into neutrinos. Contrarily to some statements given in the literature, this channel is not always suppressed, depending on the neutrino mass models. In particular for scalar mediators, it might well be possible to find annihilation rates in a detectable region. Even though the neutrino mass is practically zero for all annihilation diagrams, the mechanism responsible for its generation can very well be decisive, since different mechanisms are constrained in different ways. We have shown how to systematically search for promising situations and how some cases relate to known models. Furthermore, we have illustrated two example considerations numerically. Further advantageous cases could be modelled in the future."
    },
    "1005/1005.1263_arXiv.txt": {
        "abstract": "{ {\\it Context:}~This paper is part of a multi-species survey of line emission from the molecular gas in the circum-nuclear disk (CND) of the Seyfert 2 galaxy NGC~1068. Unlike in other active galaxies, the intensely star-forming regions in NGC~1068 and the CND can be resolved with current instrumentation. This makes this galaxy an optimal test-bed to probe the effects of AGN on the molecular medium at $\\sim100$~pc scales.  {\\it Aims:}~Single-dish observations have provided evidence that the abundance of silicon monoxide (SiO) in the CND of NGC~1068 is enhanced by 3--4 orders of magnitude with respect to the values typically measured in quiescent molecular gas in the Galaxy. % We aim at  unveiling the mechanism(s) underlying the SiO enhancement.%  {\\it Methods:}~We have imaged with the IRAM Plateau de Bure interferometer (PdBI) the emission of the SiO(2--1) (86.8~GHz) and CN(2--1) (226.8~GHz) lines in NGC~1068 at $\\sim150$~pc and $60$~pc spatial resolution, respectively. We have also obtained complementary IRAM 30m observations of HNCO and methanol (CH$_3$OH) lines. These species are known as tracers of shocks in the Galaxy.  {\\it Results:}~SiO is detected in a disk of $\\sim$400pc size around the AGN. SiO abundances in the CND of $\\sim$(1--5)$\\times$10$^{-9}$ are about 1--2 orders of magnitude above those measured in the starburst ring. % The overall abundance of CN in the CND is high: $\\sim$(0.2--1)$\\times$10$^{-7}$.  The abundances of SiO and CN are enhanced at the extreme velocities of gas associated with non-circular motions close to the AGN ($r<$70~pc). On average, HNCO/SiO and CH$_3$OH/SiO line ratios in the CND are similar to those measured in prototypical shocked regions in our Galaxy. However, the strength and abundance of CN in NGC~1068 can be explained neither by shocks nor by PDR chemistry. Abundances measured for CN and SiO, and the correlation of CN/CO and SiO/CO ratios with hard X-ray irradiation, suggest that the CND of NGC~1068 has become a giant X-ray dominated region (XDR). ",
        "introduction": "Nuclear activity and intense star formation can shape the excitation and chemistry of molecular gas in the circum-nuclear disks (CND) of galaxies.  The effect on molecular gas properties of strong radiation fields (UV and X-rays) and the injection of mechanical energy (by gas outflows or jet-ISM interactions) are key ingredients in the feedback of activity. Molecular line observations can be used to study the nature of the dominant source of energy in galaxies hosting both active galactic nuclei (AGN) and intense star formation . Furthermore, nearby galaxies can serve as local templates of distant galaxies where these phenomena can be deeply embedded.    Some of the tracers of the dense molecular gas phase, which are commonly used in extragalactic research, can be heavily affected by the feedback of AGN activity. In particular the lines of HCN (the most widely observed tracer of the dense molecular medium) appear as {\\it over-luminous} in the CND of many Seyfert galaxies with respect to other tracers of the dense molecular gas (Tacconi et al.~\\cite{Tac94}; Kohno et al.~\\cite{Koh01}; Usero et al.~\\cite{Use04}, hereafter U04; Krips et al.~\\cite{Kri08}). Doubts have been thus cast on the universality of the conversion factor between the luminosity of the HCN(1--0) line, $L_{\\rm HCN(1-0)}$, and the mass of dense molecular gas in active galaxies (Graci\\'a-Carpio et al.~\\cite{Gra06, Gra08}; Garc\\'{\\i}a-Burillo et al.~\\cite{Gar06}; Krips et al.~\\cite{Kri08}). X-rays can efficiently process large column densities of molecular gas around AGNs, producing X-ray dominated regions (XDR). It has been proposed that the abundances of certain ions, radicals and molecular species, like HCN, can be enhanced in XDR (Lepp \\& Dalgarno~\\cite{Lep96}; Maloney et al.~\\cite{Mal96}; Meijerink \\& Spaans~\\cite{Meij05}; Meijerink et al.~\\cite{Meij07}). However, it is still controversial whether pure gas-phase XDR models are able to enhance HCN abundances in AGN to the level imposed by observations.   The inclusion of dust grain chemistry, not taken into account by gas-phase XDR schemes, could solve the HCN controversy in AGN by enhancing the abundance of this molecule. Besides HCN, other molecular species can also undergo significant changes in their abundances due to dust grain processing. X-rays can evaporate small ($\\sim$10~\\AA) silicate grains (Voit~\\cite{Voi91}). This can increase the Si fraction in the gas phase and then considerably enhance the abundance of SiO in X-ray irradiated molecular gas (Mart\\'{\\i}n-Pintado et al.~\\cite{Mar00}; U04; Garc{\\'{\\i}}a-Burillo et al.~\\cite{Gar08}; Amo-Baladr\\'on et al.~\\cite{Amo09}). Furthermore, mechanical sputtering of dust grains in molecular shocks is an additional source of dust grain chemistry. Large SiO abundances have been found in the nuclei of a number of non-AGN galaxies where interferometer maps have been key to identifying large-scale molecular shocks (Garc\\'{\\i}a-Burillo et al.~\\cite{Gar00,Gar01}; Usero et al.~\\cite{Use06}).  To address the role of X-rays and shocks in shaping the chemistry of molecular gas in AGN requires a multi-species approach. Unlike SiO, other tracers of the dense molecular gas, like CN, are not expected to be highly abundant in shocks (Mitchell~\\cite{Mit84}; Fuente et al.~\\cite{Fue05}; Rodr\\'{\\i}guez-Fern\\' andez et al.~\\cite{Rod10}). However, there is a consensus supported by observations and theoretical models that the CN radical is a privileged tracer of highly ionized molecular gas, typical of photon or X-ray dominated regions (PDR or XDR) (Boger \\& Sternberg~\\cite{Bog05}; Fuente et al.~\\cite{Fue93, Fue08}; Janssen et al.~\\cite{Jan95}; Lepp \\& Dalgarno~\\cite{Lep96}; Meijerink \\& Spaans~\\cite{Meij05}; Meijerink et al.~\\cite{Meij07}).  Observations of tracers like SiO and CN can then be used to quantify the relevance of shocks and XDR chemistry in AGN.  Furthermore the high-spatial resolution provided by interferometers is paramount to spatially discriminate between the different chemical environments (star formation vs AGN activity) that can coexist in a single galaxy.   NGC\\,1068 is the strongest nearby Seyfert 2 galaxy, and as such is a prime candidate for studying the feeding and the feedback of activity using molecular line observations. Tacconi et al.~(\\cite{Tac94}) and Schinnerer et al.~(\\cite{Sch00}; hereafter S00) used the Plateau de Bure Interferometer (PdBI) to map at high-resolution (1--4$\\arcsec$) the emission of molecular gas in the central $r$\\,$\\sim$1.5-2~kpc disk using HCN and CO lines. The CO maps spatially resolve the distribution of molecular gas in the disk, showing a prominent starburst ring (hereafter SB ring) of $\\sim$1-1.5\\,kpc--radius, which contributes significantly to the total CO luminosity of NGC\\,1068 (see also Planesas et al.~\\cite{Pla91} and Baker~\\cite{Bak00}, hereafter B00). Furthermore, CO emission is detected in a central $r$\\,$\\sim$200\\,pc CND surrounding the AGN. In contrast to CO, emission from the HCN(1--0) line, which is also spatially resolved in the Tacconi et al.'s maps, arises mainly from the CND. Tacconi et al.~(\\cite{Tac94}) derived a high HCN/CO intensity ratio ($\\sim$1) in the CND. This is about a factor of 5--10 higher than the average ratio measured in the SB ring (Usero et al.~in prep). Radiative transfer calculations based on multi-line observations of HCN and CO showed that the abundance of HCN relative to CO was significantly enhanced in the CND: HCN/CO$\\sim$10$^{-3}$ (Sternberg et al.~\\cite{Ste94}).   U04 used the IRAM 30m telescope to observe with low to moderate spatial resolutions ($\\sim$10--30$\\arcsec$) the emission of eight molecular species in the CND of NGC\\,1068. The global analysis of the survey, which includes several lines of SiO, CN, HCO, H$^{13}$CO$^{+}$, H$^{12}$CO$^{+}$, HOC$^{+}$, HCN, CS, and CO, indicated that the bulk of the molecular gas emission in the CND of NGC\\,1068 could be interpreted as coming from a giant XDR. More recently, P\\'erez-Beaupuits et al.~(\\cite{Per07, Per09}) analyzed new single-dish data obtained for several rotational lines of HCO$^+$, HNC, CN and HCN, and came to similar conclusions as U04. However, the low spatial resolution of these single-dish observations are not optimal to precisely disentangle the contributions of the starburst and the CND in NGC~1068.   In this paper, we use the high-spatial resolution ($\\sim$1--3$\\arcsec$) afforded by PdBI to map the emission of the ($v$=0, $J$=2--1) line of SiO and the  $N$=2--1 transition of CN in the central $r$\\,$\\sim$1.5-2~kpc disk of NGC~1068. The spatial resolution of the new observations, an order of magnitude higher than that of the 30m survey of U04, allow us to neatly separate the emission of the SB ring from that of the CND. Furthermore, the SiO and CN PdBI maps, together with complementary single-dish data obtained in CH$_ 3$OH and HNCO lines, are used to study the chemical differentiation inside the CND through an analysis of line ratio maps. Line ratios are interpreted with the help of one-phase large velocity gradient (LVG) models. While the emission of the different lines and species analyzed in this work likely comes from  regions characterized by different physical conditions, the adopted one-phase approach allows us to explore the existence of overall trends in the chemical abundances of these species, using a minimum set of free parameters in the fit. We explore the dependence of line ratios with the illumination of molecular gas by the X-ray AGN source. We also explore the link between shock chemistry and gas kinematics in the CND.     We describe in Sect.~\\ref{Observations} the observations, including high-resolution SiO, CN and CO maps obtained with the PdBI, single-dish data obtained with the 30m telescope as well as X-ray images taken by Chandra of NGC\\,1068. Sect.~\\ref{Continuum} presents the continuum maps derived at 86.8\\,GHz and 226.8\\,GHz. The distribution and kinematics of molecular gas derived from SiO and CN are described in Sects.~\\ref{Distribution} and ~\\ref{Kinematics}.  SiO and CN abundances are discussed in Sect.~\\ref{$X$}.  In Sect.~\\ref{Models} we interpret the line ratio maps in terms of two different chemical scenarios (shocks and XDR). The main conclusions of this work are summarized in Sect.~\\ref{Summary and conclusions}. We assume a distance to NGC\\,1068 of $D\\sim$14~Mpc (Bland-Hawthorn et al.\\cite{Bla97}); this implies a spatial scale of $\\sim$70~pc/$\\arcsec$. ",
        "conclusions": ""
    },
    "1005/1005.3099_arXiv.txt": {
        "abstract": "The asymptotic response of donor stars in interacting binary systems to very rapid mass loss is characterized by adiabatic expansion throughout their interiors. In this limit, energy generation and heat flow through the stellar interior can be neglected. We model this response by constructing model sequences, beginning with a donor star filling its Roche lobe at an arbitrary point in its evolution, holding its specific entropy and composition profiles fixed as mass is removed from the surface. The stellar interior remains in hydrostatic equilibrium. Luminosity profiles in these adiabatic models of mass-losing stars can be reconstructed from the specific entropy profiles and their gradients. These approximations are validated by comparison with time-dependent binary mass transfer calculations. We describe how adiabatic mass loss sequences can be used to quantify threshold conditions for dynamical time scale mass transfer, and to establish the range of post-common envelope binaries that are allowed energetically.  In dynamical time scale mass transfer, the adiabatic response of the donor star drives it to expand beyond its Roche lobe, leading to runaway mass transfer and the formation of a common envelope with its companion star. For donor stars with surface convection zones of any significant depth, this runaway condition is encountered early in mass transfer, if at all; but for main sequence stars with radiative envelopes, it may be encountered after a prolonged phase of thermal time scale mass transfer, a so-called delayed dynamical instability. We identify the critical binary mass ratio for the onset of dynamical time scale mass transfer as that ratio for which the adiabatic response of the donor star radius to mass loss matches that of its Roche lobe at some point during mass transfer; if the ratio of donor to accretor masses exceeds this critical value, dynamical time scale mass transfer ensues.  In common envelope evolution, the dissipation of orbital energy of the binary provides the energy to eject the common envelope; the energy budget for this process consists essentially of the initial orbital energy of the binary and the initial self-energies of the binary components. We emphasize that, because stellar core and envelope contribute mutually to each other's gravitational potential energy, proper evaluation of the total energy of a star requires integration over the entire stellar interior, not the ejected envelope alone as commonly assumed. We show that the change in total energy of the donor star, as a function of its remaining mass along an adiabatic mass-loss sequence, can be calculated either by integration over initial and final models, or by a path integral along the mass-loss sequence. That change in total energy of the donor star, combined with the requirement that both remnant donor and its companion star fit within their respective Roche lobes, then circumscribes energetically possible survivors of common envelope evolution. ",
        "introduction": "\\label{intro}  Mass transfer is the signature feature of close binary evolution. In relatively wide binaries, this mass transfer can occur via capture of some part of the stellar wind of the donor by the accretor, but in close systems, the dominant mode of mass transfer is by Roche lobe overflow (RLOF). Hydrodynamic estimates of Roche lobe overflow rates \\citep{pacz72,egg06} show that mass transfer rates are typically of order \\begin{equation} \\dot{M}_1 \\approx -\\frac{M_1}{P_{\\rm orb}} \\left( \\frac{\\Delta R_1}{R_{\\rm L}} \\right)^{n + \\frac{3}{2}} \\ , \\label{Mdot} \\end{equation} where $M_1$ is the mass of the donor star, $P_{\\rm orb}$ the orbital period of the binary, $R_{\\rm L}$ the effective Roche lobe radius, which is the radius of a sphere of the same volume as the interior of the Roche lobe, $\\Delta R_1 = R_1 - R_{\\rm L}$ the donor's radius $R_1$ minus its Roche lobe radius $R_{\\rm L}$, and $n$ the donor's effective polytropic index. This mass transfer rate is exquisitely sensitive to $\\Delta R_1$; so long as mass transfer occurs on a time scale much longer than the dynamical time scale of the binary, $R_1 \\approx R_{\\rm L}$.  We can examine the stability of mass transfer by comparing the responses of donor star radius and Roche lobe radius to perturbations in the mass transfer rate. So long as $\\Delta R_1 \\ll R_{\\rm L}$, hydrostatic equilibrium prevails through the interior of the donor star, and dynamical terms become important only in the vicinity of the inner Lagrangian point. Variations in the radius of the donor can be separated into mass- and time-dependent parts: \\begin{equation} {\\rm d} \\ln R_1 = \\left( \\frac{\\partial \\ln R_1}{\\partial t} \\right)_{M_1} {\\rm d}t + \\left( \\frac{\\partial \\ln R_1}{\\partial \\ln M_1} \\right)_t {\\rm d} \\ln M_1 \\ . \\label{dlnR1} \\end{equation} The first term on the right-hand side represents not only the evolutionary expansion (or contraction) of the donor star in the absence of mass transfer, but also the expansion or contraction due to thermal relaxation if, for example, mass transfer has driven the donor out of thermal equilibrium. The second term on the right-hand side represents the hydrostatic response of the radius to changing mass, with thermal or chemical diffusion and nuclear evolution suppressed; that is, it represents the adiabatic response of the donor star to mass loss. We will henceforth label this logarithmic derivative by the subscript `ad' instead of $t$, and define \\begin{equation} \\zeta_{\\rm ad} \\equiv \\left( \\frac{\\partial \\ln R_{1}}{\\partial \\ln M_{1}} \\right)_t \\ . \\label{zetaad} \\end{equation} The Roche lobe radius of the donor depends on three quantities (assuming a circular orbit) -- its mass $M_1$, the total mass of the binary $M$, and the orbital angular momentum of the binary $J$: \\begin{equation} R_{\\rm L} = r_{\\rm L}(M_1/M) \\frac{J^2M}{G M_1^2(M - M_1)^2} \\ , \\label{radiuslobe1} \\end{equation} where $r_{\\rm L}(M_1/M)$ is the Roche lobe radius in units of the binary separation, and is a function only of the binary mass ratio, or as written here, of the ratio of donor to total mass. We can thus write \\begin{equation} {\\rm d} \\ln R_{\\rm L} = \\left( \\frac{\\partial \\ln R_{\\rm L}}{\\partial \\ln M_1} \\right)_{M,J} {\\rm d} \\ln M_1 + \\left( \\frac{\\partial \\ln R_{\\rm L}}{\\partial \\ln M} \\right)_{M_1,J} {\\rm d} \\ln M + \\left( \\frac{\\partial \\ln R_{\\rm L}}{\\partial \\ln J} \\right)_{M_1,M} {\\rm d} \\ln J \\ . \\label{dlnRL_mmj} \\end{equation} In conservative mass transfer, as that term is commonly used, $M$ and $J$ are conserved quantities, and therefore the latter two terms on the right-hand side of Eq.~(\\ref{dlnRL_mmj}) vanish. However, more generally, a binary may suffer background mass loss through a stellar wind from the companion star, or angular momentum losses through, for example, stellar winds or gravitational wave radiation. Moreover, systemic mass and angular momentum losses may contain contributions arising from a loss of part of the mass transfer stream between binary components; we can call these consequential losses. We therefore decompose the terms in ${\\rm d} \\ln M$ and ${\\rm d} \\ln J$ into terms linearly independent of ${\\rm d} \\ln M_1$ and given the subscript $M_1$, and those consequential losses linearly proportional to ${\\rm d} \\ln M_1$ and given subscript $t$, by analogy with Eq.~(\\ref{dlnR1}). We then write \\begin{equation} {\\rm d} \\ln R_{\\rm L} = \\left( \\frac{\\partial \\ln R_{\\rm L}}{\\partial t} \\right)_{M_1} {\\rm d}t + \\left( \\frac{\\partial \\ln R_{\\rm L}}{\\partial \\ln M_1} \\right)_t {\\rm d} \\ln M_1 \\ , \\label{dlnRL} \\end{equation} where \\begin{equation} \\left( \\frac{\\partial \\ln R_{\\rm L}}{\\partial t} \\right)_{M_1} = \\left( \\frac{\\partial \\ln R_{\\rm L}}{\\partial \\ln M} \\right)_{M_1,J} \\left( \\frac{\\partial \\ln M}{\\partial t} \\right)_{M_1} + \\left( \\frac{\\partial \\ln R_{\\rm L}}{\\partial \\ln J} \\right)_{M_1,M} \\left( \\frac{\\partial \\ln J}{\\partial t} \\right)_{M_1} \\label{dlnRLdt} \\end{equation} and \\begin{eqnarray} \\left( \\frac{\\partial \\ln R_{\\rm L}}{\\partial \\ln M_1} \\right)_t & \\!\\!\\! = \\!\\!\\! & \\left( \\frac{\\partial \\ln R_{\\rm L}}{\\partial \\ln M_1} \\right)_{M,J} + \\left( \\frac{\\partial \\ln R_{\\rm L}}{\\partial \\ln M} \\right)_{M_1,J} \\left( \\frac{\\partial \\ln M}{\\partial \\ln M_1} \\right)_t \\nonumber \\\\ & & \\mbox{} + \\left( \\frac{\\partial \\ln R_{\\rm L}}{\\partial \\ln J} \\right)_{M_1,M} \\left( \\frac{\\partial \\ln J}{\\partial \\ln M_1} \\right)_t , \\label{dlnRLdM} \\end{eqnarray} By analogy with Eq.~(\\ref{zetaad}) above, we define \\begin{equation} \\zeta_{\\rm L} \\equiv \\left( \\frac{\\partial \\ln R_{\\rm L}}{\\partial \\ln M_{1}} \\right)_t \\ . \\label{zetaL} \\end{equation} Combining Eqs.~(\\ref{dlnR1}) and~(\\ref{dlnRL}), we obtain \\begin{equation} {\\rm d} \\ln {R_1/R_{\\rm L}} = \\left[ \\left( \\frac{\\partial \\ln R_1}{\\partial t} \\right)_{M_1} - \\left( \\frac{\\partial \\ln R_{\\rm L}}{\\partial t} \\right)_{M_1} \\right] {\\rm d}t + \\left( \\zeta_{\\rm ad} - \\zeta_{\\rm L} \\right) {\\rm d} \\ln M_1 \\ . \\label{dlnDelR} \\end{equation}  Now suppose that the donor star exactly fills its Roche lobe. In order for mass transfer to occur, $R_1$ must increase faster (or decrease slower) than $R_{\\rm L}$ would if it were not for mass transfer; that is, the coefficient of ${\\rm d}t$ in Eq.~(\\ref{dlnDelR}) must be positive. If the mass transfer rate is self-limiting, the stellar radius will parallel the Roche lobe radius as mass transfer proceeds, i.e., ${\\rm d} \\ln (R_1/R_{\\rm L}) \\approx 0$. This condition then requires that the second term on the right-hand side of Eq.~(\\ref{dlnDelR}) be negative. Since ${\\rm d} \\ln M_1 < 0$, it follows that self-limiting (stable) mass transfer can only occur if $\\zeta_{\\rm ad} > \\zeta_{\\rm L}$. In that case, we can solve directly for the mass transfer rate: \\begin{equation} \\dot{M}_1 = - \\frac{M_1}{\\zeta_{\\rm ad} - \\zeta_{\\rm L}} \\left[ \\left( \\frac{\\partial \\ln R_1}{\\partial t} \\right)_{M_1} - \\left( \\frac{\\partial \\ln R_{\\rm L}}{\\partial t} \\right)_{M_1} \\right] \\ . \\label{zetaad_zetalobe} \\end{equation} Typically, $\\zeta_{\\rm ad}$ and $\\zeta_{\\rm L}$ are of order unity, and mass transfer is driven on a time scale $R_1/\\dot{R}_1$. If the donor remains in thermal equilibrium, mass transfer proceeds on a nuclear time scale. If, on the other hand, the donor has been driven out of thermal equilibrium by mass loss, it typically expands much more rapidly, on its thermal time scale, and accordingly mass transfer proceeds on that thermal time scale. In some circumstances, such as among the shorter-period cataclysmic variables or double-degenerate binaries, the driving time scale may be that characterizing contraction of the Roche lobe by angular momentum losses.  Dynamical time scale mass transfer occurs if $\\zeta_{\\rm ad} \\le \\zeta_{\\rm L}$. Both right-hand terms in Eq.~(\\ref{dlnDelR}) are then positive, and the donor star radius is then driven by simple hydrostatic relaxation far beyond its Roche lobe radius. In this case, the mass transfer rate can in principle accelerate up to a truly dynamical rate (cf. Eq.~(\\ref{Mdot})). Three-dimensional simulations of this process \\citep{rasi96,rick06,rick08} show that the donor envelope is tidally disrupted, engulfing the companion star while tidal torques from the lagging, distorted envelope plunge that companion and the core of the donor into orbits deeply embedded in that envelope. The binary enters common envelope evolution \\citep{pacz76}. The common envelope itself lags in rotation behind the embedded cores, which spiral inward as they lose orbital energy and angular momentum to the envelope. Ultimately, either the dissipation of orbital energy suffices to eject the common envelope, revealing the core of the original donor and its companion star\\footnote{Because the common envelope has a much higher specific entropy than the surface of the accreting companion, it is much hotter and more tenuous at that interface. The companion can accrete very little matter during common envelope evolution -- see \\citet{hjel91}.}, or the cores merge within the common envelope. Examples of the products of common envelope evolution include planetary nebulae with close binary nuclei, cataclysmic variables, and other short-period binaries containing at least one degenerate component. It is important to recognize that common envelope evolution must be very rapid, lest radiative losses quench common envelope ejection, and that energy sources other than orbital energy dissipation are unlikely to contribute significantly to the energy budget of the system \\citep{webb07}. Adiabatic mass loss sequences can therefore define the energy difference between initial and final states. This energy difference must be supplied by orbital energy dissipation\\footnote{common envelope systems must, of course, also observe angular momentum conservation \\citep[cf.][]{nele00,nele01,nele05}, but this is almost always a less stringent constraint on the final system.}.  Useful insights into the behavior of donor stars under adiabatic mass loss can be gleaned from very simplified models. \\citet{hjel87} investigated the properties of polytropic models with power-law equations of state, including complete polytropes, composite polytropes ($n = 3$ cores with $n = \\frac{3}{2}$ envelopes and $\\gamma = \\frac{5}{3}$ equations of state), and condensed polytropes ($n = \\frac{3}{2}$ envelopes with point mass cores). These models are relevant for some main-sequence stars and giant-branch stars, and illuminate the qualitative (and, within limits, quantitative) behavior of those stars. They were further employed by \\citet{sob97} to explore stability criteria for binary mass transfer. However, these models leave many important evolutionary stages of a donor unexplored, and they leave considerable room for improvement even where they may be applicable \\citep{egg89,han02,pods02,chen08}. A more sophisticated approach, using realistic stellar models, is both desirable and possible, as shown by \\citet{hjel89a,hjel89b}. This paper represents a step toward filling this need.  We describe here a method of construction of stellar adiabatic mass loss sequences. In a given sequence, the initial model corresponds to a full stellar model at an arbitrary state in its evolution. As in the pioneering work of \\citet{hjel89a,hjel89b}, mass is removed from the surface while stellar interior at each point relaxes adiabatically to a new state of hydrostatic equilibrium. In Section 2, we describe the modifications to the normal set of stellar structure equations that are needed to construct adiabatic mass loss sequences; how one can recover the internal luminosity profile, and hence thermal relaxation rate, from the adiabatic sequences; and how we have implemented these concepts in a numerical code. Adiabatic mass loss sequences constructed in this way yield $\\zeta_{\\rm ad}$, fixing the critical mass ratio for dynamical time scale mass transfer, but they also reveal thresholds for a delayed dynamical instability arising when the donor is stripped down to a nearly isentropic core. These applications are described in Section 3. In Section 4, we discuss the calculation of the total energy of a donor star as it loses mass adiabatically, and the application of those results to determining in what orbit, if any, a binary can survive common envelope evolution. We conclude with a brief appraisal of the relevance of our results to population synthesis studies in Section 5. ",
        "conclusions": "\\label{conclusions}  In this paper we have described the construction and interpretation of model sequences describing the asymptotic responses of stars to mass loss in a binary system in the adiabatic limit, in which mass transfer is so rapid that negligible heat transport occurs within the stellar interior during mass loss.  We have shown how the radial responses revealed by these adiabatic mass loss sequences can be used to establish threshold conditions for dynamical time scale mass transfer.  That instability may occur promptly upon the donor filling its Roche lobe, if the donor has a surface convection zone of any significant depth, or following an extended period of thermal time scale mass transfer (a delayed dynamical instability) when a radiative donor is stripped down to a nearly isentropic core.  Among very luminous stars with deep surface convection zones (bright giant branch and asymptotic giant branch stars), the adiabatic mass loss sequences show a pronounced initial expansion, as their superadiabatic surface layers are stripped away, complicating the interpretation of these models.  Adiabatic mass loss sequences can also be used to constrain the survival of binaries entering common envelope evolution through dynamical time scale mass transfer.  We correct certain conceptual errors appearing in the literature regarding the calculation of envelope binding energy, and describe complementary approaches to its evaluation.  By combining energy constraints with requirements that both binary components fit within their post-common envelope Roche lobes, we can place strict limits on the masses, mass ratios, and remnant orbital separations of binaries passing through common envelope evolution.  In future papers we will survey the Hertzsprung-Russell diagram, mapping out the the threshold conditions for dynamical time scale mass transfer, and for the survival of common envelope evolution, across the full range of masses and evolutionary stages of potential interest, from main sequence stars across the Hertzsprung gap, to giant branch stars and advanced stages of nuclear burning.  These are essential considerations for binary population synthesis models \\citep[e.g., ][]{han94,han95,port96,lip96,hurl02,tut02,will05,belcz08}"
    },
    "1005/1005.1864_arXiv.txt": {
        "abstract": " ",
        "introduction": "The components and properties of interiors in neutron stars have attracted much attention since the discovery of the first pulsar and then the confirmation of it to be a fast rotational neutron star. However, the composition of matter for neutron star at super-nuclear density is still an indeterminacy due to the uncertain nuclear physics. This is also the main central question for the astronomy and astrophysics decadal survey in the future~\\cite{fn}. Since the observational phenomena of pulsars and compact radioactive source with high energy are often related to the nuclear processes and the composition of dense matter in compact stars, it offers us an effective method to probe the signal of matter composition at super-nuclear density. On one hand, the equation of state (EOS) and dynamic properties would affect the structure and evolution behavior of compact stars, on the other hand, we may estimate the correlative information of nuclear matter according to the observational macroscopic compact star properties such as the structure and evolution behavior. Hence, many investigators spontaneously expect to probe the super-dense matter through astrophysical observations~\\cite{all}. However, up to now, constraints of inferred masses and radii for pulsars on equations of state cannot uniquely examine and distinguish the compositions inside neutron stars. Along with the abundance of data about spin frequency for fast rotating pulsars, researchers begin to pay much attention to confirming limit spin of pulsars, which is another way to probe into the interior of compact stars. ",
        "conclusions": "The composition of matter in the core of neutron stars has attracted much attention owing to its important significance. In our estimation, we try to probe the inner components of rapidly rotating compact stars such as the millisecond pulsar SAX J1808.4-3658 and the possible sub-millisecond pulsar XTE J1739-285 in our own way by comparing the genuine rotation frequencies under different theoretical models with the observational data, which may exert more stringent constraint on matter composition of compact stars. According to our treatment, the SAX J1808.4-3658 is a star with exotic matter and XTE J1739-285 a hybrid star."
    },
    "1005/1005.3875_arXiv.txt": {
        "abstract": "Numerical simulations, based on the core-nucleated accretion model, are presented for the formation of Jupiter at 5.2 AU in 3 primordial disks with three different assumed values of the surface density of solid particles. The grain opacities in the envelope of the protoplanet are computed using a detailed model that includes settling and coagulation of grains and that incorporates a recalculation of the grain size distribution at each point in time and space. We generally find lower opacities than the 2\\% of interstellar values used in previous calculations [Hubickyj, O., Bodenheimer, P., Lissauer, J.~J., 2005. Icarus 179, 415--431; Lissauer, J.~J., Hubickyj, O., D'Angelo, G., Bodenheimer, P., 2009. Icarus 199, 338-350].  These lower opacities result in more rapid heat loss from and more rapid contraction of the protoplanetary envelope. For a given surface density of solids, the new calculations result in a substantial speedup  in formation time as compared with those previous calculations.  Formation times are calculated to be $1.0$, $1.9$, and $4.0$ Myr, and solid core masses are found to be $16.8$, $8.9$, and $4.7\\;\\mathrm{M}_{\\oplus}$, for solid surface densities, $\\sigma$, of $10$, $6$, and $4\\;\\mathrm{g\\;cm^{-2}}$, respectively. For $\\sigma=10$ and $\\sigma=6\\;\\mathrm{g\\;cm^{-2}}$, respectively, these formation times are reduced by more than 50\\% and more than 80\\% compared with those in a previously published calculation with the old approximation to the opacity. ",
        "introduction": "} The core-nucleated  accretion theory of giant planet formation was effectively originated by \\citet{safronov72}, who realized that a major part of the process involved the accumulation of solid bodies. Consider the accretion of relatively small objects (planetesimals) onto a significantly larger embryo. If $M_{\\rm core}$ is the mass of the embryo, then the fundamental equation for its growth, in the absence of gas, is \\begin{equation} \\frac{{\\rm d}M_{\\rm core}}{{\\rm d}t} =  \\pi R_{\\rm core}^2 \\sigma\\Omega F_g \\approx \\pi R_{\\rm core}^2 \\sigma\\Omega \\bigl[1 + \\bigl(\\frac{v_e}{v}\\bigr)^2\\bigr], \\label{eq:saf} \\end{equation} where $\\pi R_{\\rm core}^2$  is the effective geometrical capture cross-section, $\\sigma$ is the surface density of solid material (planetesimals), $v_e$ is the escape velocity from the embryo, $\\Omega$ is the orbital frequency, and $v$ is the relative velocity of embryo and planetesimal. The quantity in brackets is  the gravitational enhancement factor $F_g$, here given  in the (2 + 2)-body approximation. \\citeauthor{safronov72} mentioned the later capture of gas to form giant planets, and \\citet{cameron73} stated that Jupiter could form by gas accretion once a solid core of about 10 M$_\\oplus$ had been accreted.  \\citet{perri74} then constructed an adiabatic, hydrostatic equilibrium model of a protoplanet with a core and a gaseous envelope, extending outward to the Hill radius.  They were able to show that there was a critical mass for the core, above which the envelope was unstable, implying that rapid gas accretion could occur. However, the value of this critical core mass was at least 70 M$_\\oplus$, much higher than  the deduced core masses of Jupiter and Saturn. The assumption of adiabaticity turns out to be incorrect \\citep{stevenson82}, as more detailed calculations show that radiative transport with subadiabatic gradients is able to carry the energy outwards in significant portions of a growing giant planet.  \\citet{mizuno80} constructed models of a protoplanet including  both radiative and convective energy transport in the envelope, along with an energy source provided by the planetesimals accreting onto the core at a constant rate, typically  $\\dot M_{\\rm core} =10^{-6}$ M$_\\oplus$/yr. The model assumed strict hydrostatic and thermal equilibrium, and the outer radius was assumed to be the protoplanet's Hill radius. Again a critical core mass was found, above which the envelope was unstable, and the value was about 12 M$_\\oplus$, close to the estimated core masses of the giant planets. Furthermore, the critical value was practically independent of the distance to the Sun, partly as a result of the assumption of the same constant value of $\\dot M_{\\rm core}$ at all distances.  The assumption of strict hydrostatic equilibrium was relaxed by \\citet{bodenheimer86}, who allowed quasi-static contraction and calculated the first evolutionary sequences applied to the core-nucleated accretion model. The planetesimal accretion rate was assumed to be constant, also at $10^{-6}$ M$_\\oplus$/yr. They found that gravitational contraction  (but not dynamical collapse) began to become important at about the point where core and envelope masses were equal, at a value which later came to be known as the {\\it crossover} mass. Rapid gas accretion also began near the crossover mass, which they found to be about 16.5 M$_\\oplus$.  The next major step was a series of calculations by \\citet{pollack96} who allowed a variable core accretion rate given by eq.~(\\ref{eq:saf}). The simulations contained three major elements: (1) three-body accretion cross-sections of solids onto the embryo \\citep{greenzweig92} to determine $F_g$, (2) a modified stellar evolution code \\citep{henyey64} to follow the structure and evolution of the gaseous envelope, and (3) a calculation of the trajectories of planetesimals through the gaseous envelope \\citep{podolak88}, to deduce the deposition of solid mass and energy in the envelope and to determine the protoplanet's effective capture radius, taking into account gas drag. The often-quoted result of these calculations is that it took Jupiter 8 Myr to form in a solar nebula with $\\sigma=10$ g cm$^{-2}$, three times the density in the minimum mass solar nebula, under the assumption that the opacity due to small grains in the cool outer envelope was provided by grains with an interstellar size distribution and with interstellar abundances with respect to the gas. It was also shown that if the grain opacities were arbitrarily reduced to 2\\% of interstellar values, the formation time was reduced to about 3 Myr.   Subsequent observations of protoplanetary disks and modelling of the interior of Jupiter revealed two major problems with the \\citet{pollack96} baseline-case results for the formation of Jupiter.  First, that   case gave a formation time longer than the mean lifetime of protostellar disks, and, second, it produced a core mass of about  16  M$_\\oplus$, above the plausible range at the time of  (0--11) M$_\\oplus$, derived from matching interior models of the present Jupiter with observations \\citep{saumon04}. Further calculations \\citep{hubickyj05} considered, with generally improved physics, the effects of the following parameters: (1) the opacity arising from small grains in the envelope, (2) the solid surface density $\\sigma$ in the protoplanetary disk, and (3) a possible cutoff in core accretion rate as a result of competition for solid particles by neighboring embryos. The reduction in opacity was justified by calculations \\citep{podolak03} showing that grains entering into the protoplanetary envelope would rapidly coagulate, settle, and eventually evaporate, resulting in an actual opacity well below interstellar values. In the actual protoplanetary formation calculation the grain opacity was simply reduced to 2\\% of interstellar values, with the result that the formation time was reduced to under     3 Myr, half that found for interstellar values of grain opacity. The final core mass, however, still remained near 16 M$_\\oplus$. A reduction in $\\sigma$ from 10 to 6 g cm$^{-2}$, again with the opacity at 2\\% of interstellar, increased the formation time to 13 Myr and reduced the core mass to 8 M$_\\oplus$. It turned out that the constraints of short formation time and small core mass could be satisfied simultaneously only if the accretion rate of solids into the protoplanet were arbitrarily cut off at some time, simulating roughly the effect of planetesimal accretion by neighboring protoplanets. A cutoff at $M_{\\rm core} = 5$ M$_\\oplus$ led to a formation time of 4.5 Myr, while a cutoff at 10 M$_\\oplus$ gave a formation time of 1 Myr.  In the planet formation simulations of \\citet{pollack96} and \\citet{hubickyj05},  simplified surface boundary conditions were applied during the rapid gas accretion phase. First, the gas density in the protoplanetary disk was assumed to be constant in time. Second, the gas accretion rate was capped at $10^{-2}$ M$_\\oplus$/yr, a rough estimate of the rate at which the disk could supply gas to the protoplanet as a result of its viscous evolution. Third, the outer boundary of the protoplanet was set at the modified accretion radius $R_A$, defined by \\citet{bodenheimer00} to be \\begin{equation} R_A = \\frac{GM_p}{c_s^2/k_1 + \\frac{GM_p}{(k_2R_H)}}, \\label{eq:ra} \\end{equation} where $R_H$ is the Hill sphere radius, $c_s$ is  the sound speed in the disk, $M_p$ is the planet's mass, and $k_1,~k_2$ are constants, both set to unity. Equation~(\\ref{eq:ra})  is an interpolation between the Bondi accretion radius and the Hill sphere radius, such that when $R_H$ is large, $R_A$ reduces to the Bondi radius, and when $R_H$ is small, $R_A$ reduces to the Hill radius.  These assumptions were modified by \\citet{lissauer09} to take into account the results of new three-dimensional numerical simulations of protoplanetary disks accreting onto protoplanets. These simulations employed the  code developed by \\citet{dangelo03}.  The results from the 3-D simulations showed that $k_2 \\approx 0.25$, that is, the planet can collect material from the disk only within a quarter of its Hill radius, at best.  The gas accretion rates from the 3-D simulations, which included the effects of gap opening, were used whenever they were less than the accretion rate needed to maintain the boundary of the planet at $R_A$; thus the arbitrary cap on this rate was removed. The parameters investigated in this set of simulations included (1) the size of the region from which the protoplanet can accrete, (2) the viscosity in the protoplanetary disk, which has a strong influence on the disk accretion rate onto the protoplanet once the envelope has shrunk and the planet is massive enough to begin clearing gas from nearby its orbit, and (3)  the time-dependence of the gas density in the protoplanetary disk in the vicinity of the protoplanet. The grain opacity was set to 2\\% interstellar and $\\sigma=10$ g~cm$^{-2}$. The results of these improved simulations showed that Jupiter's formation time was still less than 3 Myr, and that a relatively low-viscosity disk ($\\alpha \\approx 4 \\times 10^{-4}$) is consistent with the formation of planets of about 1 Jovian mass, while higher viscosity produced planets that were more massive than most observed extrasolar giant planets.  In all of the simulations just mentioned, as well as in detailed formation simulations by other groups \\citep{alibert05,fortier07,ikoma00}, the opacities arising from grains in the protostellar envelope have been approximated either by interstellar grain values \\citep{pollack85} or by a reduction of those values by some arbitrary amount, typically a factor 50. However, the simulations of \\citet{podolak03} and \\citet{movshovitz08}, in which grain coagulation and sedimentation were calculated for a few particular protoplanetary envelopes, showed that the ratio of actual grain opacity to interstellar grain opacity was highly variable, depending on the depth in the envelope.  Near the surface of the planet the opacities were close to interstellar, while deep in the envelope they were reduced by a factor of up to 1000. Those  opacity calculations, however, were  not coupled to the evolution of the protoplanet, and they included only the outer radiative layer in the atmosphere, which typically has a temperature of 600 K at its inner boundary.  In the present paper we extend the basic grain settling and coagulation model of  \\citet{movshovitz08} by including  the entire depth of the envelope down to a temperature of 1800 K, where practically all grains have evaporated. Since some layers of the outer part of the envelope are convective, we also modify the grain simulation to take into account convective effects. We then couple the grain calculation to the evolution of the protoplanet, recalculating the grain size distribution and opacity in every layer of the protoplanet at every time step of the evolutionary calculation. These significant improvements in the core accretion model allow us to investigate the following questions: (1) How is the formation time for a Jupiter-mass planet  at 5.2 AU modified compared with the earlier approximate calculations regarding the opacity? (2)  Most of our previous simulations for the formation of Jupiter required an initial solid surface density of 10 g cm$^{-2}$ in the protoplanetary disk. How far can this assumed density be reduced such that Jupiter still forms within a dissipation time of the disk -- roughly 3 Myr -- and what is the resulting solid core mass? The following sections describe the nature of the grain simulations, the basic idea of the core-nucleated  accretion model, how the two calculations are integrated, and our new results for four formation runs, up to the onset of rapid gas accretion, for a Jupiter-mass planet at 5.2 AU. ",
        "conclusions": "The formation phase of Jupiter, as determined by the concurrent accretion of solid particles and gas, has been calculated under the following major assumptions: (1) The protoplanet remains fixed  in position at 5.2 AU. (2) The planet is assumed to exhibit runaway accretion of solids within a disk of 100 km planetesimals. The planetesimals are well mixed within the feeding zone of the planet; thus their surface density is constant in space but varies in time as the planet accretes. Migration of planetesimals into or out of the feeding zone is not considered. (3) The protoplanet is assumed to be the sole dominant mass in the region; other embryos which might compete for planetesimals are not considered. (4) The solids accretion rate is determined from Safronov's formula (eq.~1) with gravitational enhancement factor obtained from the 3-body numerical fits of \\citet{greenzweig92}.  (5) All of the planetesimal material accreted by the protoplanet is assumed to fall into the core; thus what is referred to as `core mass' in this paper is really the total heavy element excess over solar abundances. (6) The gas accretion rate is determined by the requirement that the outer boundary of the protoplanet remain at the modified accretion radius (eq.~2). Three-dimensional simulations of a planet embedded in a disk \\citep{lissauer09} determine in part the region of a planet's gravitational influence.  (7) The opacity caused by grains in the envelope of the protoplanet is determined by a detailed calculation of grain settling and coagulation, with the grain size distribution and the resulting grain opacity being recalculated at every level in the envelope at every time step. It is this last assumption that represents a major improvement over past calculations of Jupiter's formation, and it results in a significant reduction in the formation time scale as compared to past calculations. There are approximations involved in the treatment of convection zones in the grain calculations, but the comparison of cases $\\sigma$10 and $\\sigma$10R shows that these have a negligible effect on the results. The grains in the incoming gas are composed of silicates only, with an abundance by mass of 1\\% that of the gas. Molecular opacities are included \\citep{freedman08}.  The main parameter varied in this set of calculations is the solid surface density $\\sigma$ at 5.2 AU in the primordial disk; values considered are 10, 6, and 4 g cm$^{-2}$. As in previous simulations of Jupiter's formation using the prescription of \\citet{lissauer87} for the runaway accretion of solids, the results show that there are three main phases. Phase 1 involves primarily the accretion of solid particles to build up the core; the mass of the gaseous envelope is fairly negligible. Phase 2 begins when the isolation mass is reached; accretion of gas takes place at about three times the rate of accretion of solid particles, and the envelope builds up in mass relative to the core. Phase 3, rapid gas accretion, begins close to the point where core and envelope masses are equal. The calculations in this paper are taken beyond the  crossover mass up to the point where the gas accretion rate becomes limited by the rate at which the disk can supply the mass, as determined by three-dimensional simulations ($1.5-3 \\times 10^{-2}$ M$_\\oplus$ yr$^{-1}$). The accretion up to the final mass of Jupiter takes place quickly after that point \\citep{lissauer09}.  The formation time to crossover is dominated by Phase 2 in all simulations. Case $\\sigma$10 may be compared with a very similar calculation \\citep{lissauer09} in which the grain opacity was simply assumed to be 2\\% that of interstellar grains. Phase 1 in the present case was completed in about the same time as in the comparison case, but Phase 2 (0.52 Myr) was a factor of 3.6 shorter, giving a total formation time of about 1 Myr, compared with 2.6 Myr in the comparison case.  Correspondingly, the mean luminosity in phase 2 was a factor of 3.6 higher in the present run than in the comparison case. The main cause of the increase in luminosity is the very low opacity in the region around 600 K in the present model (Fig.~\\ref{fig:f3}), with a minimum of opacity of about $6 \\times 10^{-4}$ cm$^2$ g$^{-1}$ as compared with $0.04$ cm$^2$ g$^{-1}$ at the same temperature in the comparison case. Some layers in the present case have opacities higher than 2\\% interstellar, but they are mainly in the outermost layers that are either convective or rather optically thin, so they do little to influence the flow of radiation. In spite of the difference in formation time for the two cases, the core mass at onset of limiting gas accretion was practically the same in both cases, about 17 M$_\\oplus$.  This mass is higher than that usually derived for the present Jupiter \\citep{saumon04} but it is close to that predicted from an independent calculation of the equation of state by \\citet{militzer08}.  The question then arises whether a reduction in $\\sigma$, which is known to result in a smaller core mass, still leads to a reasonable formation time. Cases $\\sigma$6 and $\\sigma$4 show that this is indeed the case. The formation times are 1.9 and 4.0 Myr, respectively. The ratios of Phase 2 times for the three cases scale as $M_{\\rm iso}^{5/3}$/$L$, as derived by \\citet{pollack96} under the assumption that planetesimal accretion provides most of the luminosity during Phase 2. The value 4 Myr is  within the range of expected lifetimes of observed circumstellar disks around young stars \\citep{haisch01}. Note that the $\\sigma6$ case has a formation time only 14\\% as long (and Phase 2 time only 8\\% as long) as the corresponding case in \\citet{hubickyj05} with grain opacities set at 2\\%  of interstellar values, while the $\\sigma10$ case took 40\\% as long as the corresponding case ($10\\mathrm{L}\\infty$) in that earlier work.   The main conclusions of this work are: (1) Within the basic framework of a core-nucleated accretion model for Jupiter's formation, the effects of grain settling and coagulation in the gaseous envelope of a protoplanet are significant in determining the gas accretion rate and formation time. (2) The formation time for a Jupiter at 5.2 AU with a solid surface density of 10 g cm$^{-2}$ (about 3 times that in the minimum mass solar nebula) is about 1 Myr, significantly shorter than that in comparable comparison cases. If $\\sigma$ is taken to be 4 g cm$^{-2}$, the formation time is still reasonably short, about 4 Myr. (3) In the case with $\\sigma =  4$ g cm$^{-2}$, Jupiter's core mass is calculated to be 4.74 M$_\\oplus$, consistent with the Jupiter structure models computed by \\citet{saumon04}. No arbitrary cutoff in the core accretion rate was needed to obtain this result.  Further work is needed, however, to confirm this result. The main points which need to be improved include, first, the planetesimal accretion model, which should be replaced by a statistical calculation taking into account the evolution of the  planetesimal size distribution and the effects of the protoplanet itself on the planetesimal velocity distribution \\citep{weidenschilling97}. Numerous estimates of the Phase 1 accretion time have been published. One example is the (statistical + N-body) calculation of \\citet{inaba03} which takes into account planetesimal fragmentation as  well as the effect of the gaseous envelope in increasing the effective cross-section. A planetesimal size of 10 km is assumed. If the opacity in the planetary envelope is reduced by a factor 100 compared with interstellar values,   a core approaching 10 M$_\\oplus$ is reached in 5 Myr at 5.2 AU; however $\\sigma$ is required to be about 5 times that in the minimum mass solar nebula. This and other simulations, along with the present results, suggest  that the formation time for Jupiter may be controlled by Phase 1, not Phase 2. Second, the effects of the ices should be included in the calculations of grain  settling and coagulation. Third, the effects of the dissolved accreted solid material, primarily ices, should be included in the equation of state and opacity of the envelope and in the calculation of the solid core mass."
    },
    "1005/1005.4980_arXiv.txt": {
        "abstract": "The Athena MHD code has been extended to integrates the motion of particles coupled with the gas via aerodynamic drag, in order to study the dynamics of gas and solids in protoplanetary disks and the formation of planetesimals. Our particle-gas hybrid scheme is based on a second order predictor-corrector method. Careful treatment of the momentum feedback on the gas guarantees exact conservation. The hybrid scheme is stable and convergent in most regimes relevant to protoplanetary disks. We describe a semi-implicit integrator generalized from the leap-frog approach. In the absence of drag force, it preserves the geometric properties of a particle orbit. We also present a fully-implicit integrator that is unconditionally stable for all regimes of particle-gas coupling. Using our hybrid code, we study the numerical convergence of the non-linear saturated state of the streaming instability. We find that gas flow properties are well converged with modest grid resolution (128 cells per pressure length $\\eta r$ for dimensionless stopping time $\\tau_s=0.1$), and equal number of particles and grid cells. On the other hand, particle clumping properties converge only at higher resolutions, and finer resolution leads to stronger clumping before convergence is reached. Finally, we find that measurement of particle transport properties resulted from the streaming instability may be subject to error of about $\\pm20\\%$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.1870_arXiv.txt": {
        "abstract": "We use dark matter haloes identified in the {\\sc MareNostrum Universe} and galaxy groups identified in the Sloan Data Release 7 galaxy catalogue, to study the relation between halo shape and halo dynamics, parametrizing out the mass of the systems. A strong shape-dynamics, independent of mass, correlation is present in the simulation data, which we find it to be due to different halo formation times. Early formation time haloes are, at the present epoch, more spherical and have higher velocity dispersions than late forming-time haloes. The halo shape-dynamics correlation, albeit weaker, survives the projection in 2D (ie., among projected shape and 1-D velocity dispersion). A similar shape-dynamics correlation, independent of mass, is also found in the SDSS DR7 groups of galaxies and in order to investigate its cause we have tested and used, as a proxy of the group formation time, a concentration parameter. We have found, as in the case of the simulated haloes, that less concentrated groups, corresponding to late formation times, have lower velocity dispersions and higher elongations than groups with higher values of concentration, corresponding to early formation times. ",
        "introduction": "\\label{intro}   According to the hierarchical model of structure formation, groups and clusters of galaxies, embedded in dark matter haloes, emerge from Gaussian primordial density fluctuations and grow by accreting smaller structures, formed earlier, along anisotropic directions. Such structures constitute therefore an important step in the hierarchy of cosmic structure formation and are extremely important in deciphering the processes of galaxy and structure formation. It should thus be expected that recently formed structures are highly elongated, reflecting exactly the anisotropic distribution of the large-scale structure from which they have accreted their mass, and dynamically young. Therefore the determination of the dynamical state of groups and clusters of galaxies and its evolution, which could be affected by a multitude of factors - like environment, formation time, etc -  is a fundamental step in investigating the hierarchical galaxy formation theories.  Both analytical calculations and N-body simulations have consistently shown that the virialization process will tend to sphericalize the initial anisotropic distribution of matter, and therefore the shape of haloes is an indication of their evolutionary stage. Furthermore, the halo shape, size, and velocity dispersion (among others) are important factors in determining halo member orbits and interaction rates, which are instrumental in understanding galaxy formation and evolution processes (eg., Jing \\& Suto 2002; Kasun \\& Evrard 2005; Avila-Reese et al. 2005; Bailin \\& Steinmetz 2005; Allgood et al. 2006; Bett et al. 2007; Ragone-Figueroa \\& Plionis 2007; Macci\\'o, Dutton, van den Bosch 2008 and references therein). It has been found that dark-matter haloes are triaxial but with a strongly preferred prolatness (eg., Plionis, Basilakos \\& Ragone-Figueroa 2006; Gottl\\\"ober \\& Yepes 2007 and references therein) and that there is also a correlation between the orientation of their major axes and the surrounding structures. Such alignment effects, among relatively massive haloes, have been found to be particularly strong (eg., Faltenbacher et al. 2002; Kasun \\& Evrard 2005; Hopkins, Bachall \\& Bode 2005; Ragone-Figueroa \\& Plionis 2007; Faltenbacher et al. 2008 and references therein). There are indications that the mentioned correlation might arise, or strengthen, from re-arrangements of the orientation of the halo axes in the direction from which the last major merger event occurred \\cite{van}. Hence, the shape of dark matter haloes might well be correlated with the large scale structure, providing clues to the fashion in which mass is aggregated (eg., Basilakos et al. 2006). The imprints of this accretion can be observed in the substructure features present in a halo, which have been found to correlate with the halo shape (Ragone-Figueroa \\& Plionis 2007). Haloes with higher levels of substructure, or equivalently dynamically young haloes, are preferentially more elongated. The subsequent evolution should induce a change in the shape of haloes, leading to more spherical configurations, increasing their velocity dispersions as well.  According to this last statement, Faltenbacher \\& Mathews (2007) have found, in agreement with the Jeans equation, that the velocity dispersion of sub-haloes increases with the host halo concentration. Given that halo concentration and formation time are also correlated (Bullock et al. 2001; Wechsler et al. 2002), this result implies that older haloes should have higher sub-halo velocity dispersions. Once a halo is virialized, the dynamical mass can be computed from their velocity dispersion and virial radius (eg., Binney \\& Tremaine 1987). Nevertheless, if virialization has not yet been achieved, the above mentioned effect and its possible dependence to the local halo environment should be taken into account. In addition, there is also evidence for the sphericalization of DM haloes with time (e.g. Allgood et al. 2006). This behavior can be well studied using haloes in numerical simulations, where their 3D shape and velocity dispersion can be calculated accurately.  These issues motivated us to study the relation between the DM halo velocity dispersion and shape, in narrow halo mass ranges, in order to exclude the possible halo-mass dependent effects. We also investigate how this correlation is transformed when computing projected shapes and velocity dispersion along the line-of-sight, in order to be able to compare with observational data, like the SDSS galaxy groups.  The plan of the paper is the following. In section \\ref{datos} we describe the {\\sc MareNostrum} numerical simulation and the SDSS observational data together with the identification procedure of the corresponding systems and their properties. In section \\ref{hsdr} we study the halo shape-dynamics relation in the simulation. In section \\ref{proxies} we seek for a formation-age proxy which can be computed in a realistic group catalog. We continue in section \\ref{2d} by presenting the shape-dynamics correlation results for the observed SDSS groups of galaxies. Finally, section \\ref{conc} contains our conclusions. ",
        "conclusions": "\\label{conc}  We have identified dark matter haloes in the {\\sc MareNostrum} Universe and groups of galaxies in the SDSS DR7 with the purpose of searching for possible, independnet of mass, dependencies of halo/galaxy groups shape on their dynamics.  We have computed the halo formation time as the redshift at which the halo accretes half of its final mass. Since this definition of formation time is not applicable to observed groups identified in realistic galaxy catalogues, we have identified a concentration parameter that can be used as a proxy of the formation time of haloes and at the same time it is easily computed for the observed SDSS galaxy systems.  On this regard we have found a significant correlation between halo shape and dynamics (velocity dispersion) which is independent of the halo mass. We have identified the cause of this correlation to be the halo formation time differences. Early formation-time haloes show a very weak such correlation, being more spherical, having a higher velocity dispersion and showing substantial concentration. On the contrary, late formation-time haloes show a strong shape-dynamics correlation (having a wide range of shape and velocity dispersion values), with typically lower velocity dispersions, lower sphericities and significant lower concentrations. We have also studied the influence of multi-core haloes, i.e., major merger systems which the FOF joins into a single system and which other halo-finders would have probably singled out each component, in the strength of the LFT shape-dynamics correlation and verified that although they enhance the correlation, it however persists even excluding this subsample of haloes.  Finally, we find a fraction of the LFT haloes showing a somewhat unexpected behaviour, having high sphericities and velocity dispersions (comparable or even higher than those of more virialized systems), a fact which could be attributed to either some late formation-time haloes having transitory spherical shapes and large velocity dispersion, during the minimum pericentric passage of a merging process, or/and to a relatively faster virialization process occuring in some of these haloes.  The halo morphology-velocity dispersion correlation, albeit weaker, survives the 2D projection, indicating that an analogous behaviour should be expected in observational group samples. Applying a similar analysis to a ``volume-limited'' subsample of the SDSS DR7 groups of galaxies, identified using a friends of friends algorithm, we have also found a group shape and dynamics correlation, which is independent of group mass (see also Tovmassian \\& Plionis 2009). Using a concentration parameter as a proxy of the formation time, we indeed find a significant, independent of mass, shape-dynamics correlation, in a similar fashion as in the dark matter halo case. Less concentrated (late forming) groups are typically less spherical and have lower velocity dispersion than their equal mass more concentrated (early forming) counterparts."
    },
    "1005/1005.1277_arXiv.txt": {
        "abstract": "Beauty and charm $e^+e^-$ factories running at resonance thresholds have unique capabilities for studies of the production of light Dark Matter particles in the decays of $B_q (D)$ meson pairs. We provide a comprehensive study of light Dark Matter production in heavy meson decays with missing energy $ \\dslash E$ in the final state, such as $B_q (D^0) \\to\\dslash E$ and $B_q (D^0) \\to \\gamma \\dslash E$. We argue that such transitions can be studied at the current flavor factories (and future super-flavor factories) by tagging the missing-energy decays with $B(D^0)$ decays ``on the other side.\" ",
        "introduction": "The presence of cold Dark Matter (DM) in our universe provides the most natural explanation for several observational puzzles, from the original measurement of the rotational curves~\\cite{Zwicky:1933gu} of galaxies to the observation of background objects in the Bullet Cluster~\\cite{Clowe:2006eq} and spectrum features of the cosmic microwave background (CMB) fluctuations. In the conventional picture, DM accounts for the majority of mass in our Universe. However, the nature of DM is still very much a mystery, which could intimately connect astronomical observations with predictions of various elementary particle theories. Many such theories, with the notable exception of the Standard Model (SM), predict one or more stable, electrically-neutral particles in their spectrum~\\cite{PrimBH}. These particles could form all or part of the non-baryonic Dark Matter in the Universe.  Different models provide different assignments for DM particles' spin and various windows for their masses and couplings to luminous matter. In the most popular models DM is a weakly interacting particle particle with mass set around the electroweak energy scale. This follows from the experimental measurements of the relic abundance $\\Omega_{DM} h^2 \\sim 0.12$ by WMAP collaboration~\\cite{Spergel:2003cb} \\begin{equation}\\label{RelAb} \\Omega_{DM} h^2 \\sim \\langle \\sigma_{ann} v_{rel}\\rangle^{-1} \\propto \\frac{M^2}{g^4} \\sim 0.12, \\end{equation} where $M$ and $g$ are the mass and the interaction strength associated with DM annihilation respectively. As one can see, a weakly-interacting massive particle (WIMP) with electroweak-scale mass naturally gives the result of Eq.~(\\ref{RelAb}). This, coupled with an observation that very light DM particles might overclose the Universe  (known as the Lee-Weinberg limit~\\cite{LWL}) , seems to exclude the possibility of the light-mass solution for DM, setting $M_{DM} > 2-6$~GeV.  A detailed look at this argument reveals that those constraints could be easily avoided, so even MeV-scale particles can be good DM candidates. For instance, DM could be non-fermionic~\\cite{Bird:2006jd,Burgess:2000yq}, in which case the usual suppression of the DM annihilation cross-section used in setting the Lee-Weinberg limit does not hold. In addition, low energy resonances could enhance the cross-section without the need for a large coupling constant. Other solutions, which also provide low-mass candidates for DM particles, are also possible~\\cite{Feng:2008ya,Pospelov:2007mp,Kim:2009ke}.  There are many experiments designed to search for both direct interactions of DM with the detector and indirect evidence of DM annihilations in our or other galaxies by looking for the products such as gamma-rays, positrons and antiprotons. Those can in principle probe low-mass DM. However, direct searches, performed by experiments such as DAMA and CDMS~\\cite{DAMA}, rely on the measurement of the kinematic recoil of the nuclei in DM interactions. For cold DM particles, such measurements lose sensitivity with the decreasing mass of the WIMP as recoil energy becomes smaller~\\cite{Petriello:2008jj}. Indirect experiments, such as HESS~\\cite{HESS}, are specifically tuned to see large energy secondaries, only possible for weak-scale WIMPs. The backgrounds for positron and antiproton searches by HEAT and/or PAMELA experiments~\\cite{PAMELA} could be prohibitively large at small energies.  It is well-known that the existing $e^+e^-$ flavor factories and future super-flavor factories could provide the perfect opportunity to search for rare processes, especially the ones that require high purity of the final states. In particular, probes of rare B-decays, such as $B \\to K^{(*)} \\nu \\overline{\\nu}$, are only possible at those machines. These colliders, where $B_q(D)$ and $\\overline{B}_q(\\overline{D})$ are produced in charge and CP-correlated states, have an opportunity to tag the decaying heavy meson ``on the other side,\" which provides the charge or CP-identification of the decaying ``signal\" $B$ or $D$ meson. In fact, many CP-violating parameters at B-factories have been measured using this method~\\cite{Antonelli:2009ws}. It is then possible to perform a similar tag on the meson decaying to a pair of light DM particles or a pair of DM particles and a photon. The latter process might become important for some DM models as it eliminates helicity suppression of the final state\\footnote{ This is similar to the situation in leptonic decays of $B$-mesons, where the branching ratios ${\\cal B}(B \\to \\mu\\nu\\gamma) \\approx {\\cal B}(B \\to \\mu\\nu)$ and ${\\cal B}(B \\to e\\nu\\gamma) \\gg {\\cal B}(B \\to e\\nu)$. }. Moreover, compared to $B \\to K + \\displaystyle{\\not} E$ transitions, where $\\displaystyle{\\not} E$ is missing energy, a massless photon could provide better experimental opportunities for tagging without reducing the probed parameter space of the DM masses. Finally, searches for light DM in heavy meson decays could be more sensitive than direct detection and other experiments, as DM couplings to heavy quarks could be enhanced, as for example happens in Higgs portal models~\\cite{HiggsPortal}.  In this paper we compute branching ratios for the heavy meson states decaying into $\\chi_s \\overline\\chi_s$ and $\\chi_s \\overline\\chi_s\\gamma$. Here $\\chi_s$ is a DM particle of spin $s$, which appears as missing energy in a detector. The DM {\\it anti-particle} $\\overline \\chi_s$ may or may not coincide with $\\chi_s$. We shall first consider model-independent interactions of DM particles of spin-0, spin-1/2, and spin-1 with quarks. In each case we write the most general effective Hamiltonian coupling DM particles to flavor-changing $b\\to q$ (where $q=s(d)$) or $c \\to u$ current and compute $B(D)\\to \\chi_s \\overline\\chi_s (\\gamma)$ decay rates. We then consider popular models, already available in the literature, that can generate those processes. ",
        "conclusions": "We have argued that missing energy decays of the heavy mesons - $B_d$, $B_s$ and $D^0$ - provide an important way to probe different properties of Dark Matter. Consideration of different decay modes - two body decays, radiative and light meson + DM decays -  restricts different regions of the Dark Matter parameter space. Combined constraints obtained from different decay modes of various heavy mesons provide indispensable probe of physics beyond the Standard Model in general and the nature of the Dark Matter in particular. For instance, observation of $B_q (D^0) \\to \\gamma \\dslash E$, but non-observation of $B_q (D^0)\\to\\dslash E$ transitions directly point to non-self-conjugated nature of scalar DM.  We reported general constraints on the Wilson coefficients of the effective operators describing interactions of DM with quarks (see Tables I-III). Restrictions obtained in our paper are much stricter than constraints from single decay modes. Our results combined with constraints from astrophysical observables (for example \\cite{Badin}), direct detection of Dark Matter and invisible decays of heavy hadrons \\cite{Gagik} could provide a full set of tools needed to test (or rule out) the models of light Dark Matter."
    },
    "1005/1005.3933_arXiv.txt": {
        "abstract": "Particle creation at the expense of gravitational field might be sufficient to explain the cosmic evolution history, without the need of dark energy at all. This phenomena has been investigated in a recent work by Lima et-al \\cite{l:s} assuming particle creation at the cost of gravitational energy in the late Universe. However, the model does not satisfy the WMAP constraint on matter-radiation equality \\cite{g:r}. Here, we have suggested a model, in the same framework, which fits perfectly with SNIa data at low redshift as well as early Integrated Sachs-Wolfe effect on matter-radiation equality determined by WMAP at high redshift. Such a model, requires the presence of nearly $26\\%$ primeval matter in the form of baryons and CDM. ",
        "introduction": "Recently released 7-year WMAP data \\cite{w1},\\cite{w2} has found no trace of deviation from the standard $\\Lambda$CDM model. But the problem in connection with the cosmological constant remains unresolved. The vacuum energy density, as calculated by the field theorists, is some $10^{120}$ order of magnitude greater than the cosmological constant $\\Lambda$ required by the cosmologists, to explain late time cosmic acceleration, which is of the order of $H_{0}^2$, $H_0$ being the present Hubble parameter. So far, many alternatives to the standard $\\Lambda$CDM model have been proposed and explored and as a matter of fact, all of these models have been found suitable to explain late time cosmic acceleration. The problems associated with these models are, - they can not be distinguished from the standard $\\Lambda$CDM model at one hand, and most of them are not suitable to explain the early Universe, on the other. However, $\\Lambda$CDM model requires $26\\%$ of matter in the form of pressureless dust, out of which only $4\\%$ are baryons and the rest, about $22\\%$ are cold dark matter (CDM). Since dark matter interacts only with the gravitational field, so  it plays a key role in the structure formation. The gravitational Jeans instability allows compact structures to form and is not opposed by any force like radiation pressure in the case of dark matter. As a result, dark matter begins to collapse into a complex network of dark matter Halos, well before ordinary baryonic matter, which is impeded by pressure force. Without dark matter the epoch of Galaxy formation would have occurred at a substantially later stage, than observed. Thus, the amount of CDM ($22\\%$) must have been created in the very early Universe, prior to the radiation dominated era together with the baryons, by some sort of mechanism, viz., supersymmetry breaking, cosmic string decay or particle creation at the expense of gravitational field. These particles are usually supposed to be weakly interacting massive particles (WIMP). For example, as a heavy stable particle, the lightest neutralino is an excellent candidate to comprise the Universe's cold dark matter. In many models (see \\cite{br} for a nice review) the lightest neutralino can be produced in the hot early Universe and leave approximately the right relic abundance to account for the observed dark matter, ie., $22\\%$ as required by $\\Lambda$CDM model. Now, if phenomenologically one considers that CDM may also be produced by gravitational particle creation mechanism, even at a very slow rate, during the late time evolution of the Universe, viz., during the matter dominated era, as considered by Lima et-al \\cite{l:s}, then it may be possible to explain the presently observable acceleration of the Universe, without taking dark energy into account. The very advantage of the creation of cold dark matter over dark energy is that it avoids coincidence problem and also may be detectable in future experiments. In the present work our focus is on the cosmological consequences of particle production on the evolution of late stage of the Universe which was initiated recently by Alcaniz and Lima \\cite{j} and Lima, Silva and Santos(lss) \\cite{l:s}. \\\\  \\noindent We remember that in the eighty's, the motivation that initiated to go after inflationary scenario, was the fundamental three problems associated with Friedmann model, viz., the horizon, the flatness and the observed isotropy and homogeneity in the cosmic microwave background radiation (CMBR). There was an additional problem in the form of huge entropy per baryon in the observable Universe $(\\sim 10^{87})$. Einstein's equations are purely adiabatic and reversible. Consequently, these equations can hardly provide, by themselves, an explanation relating to the origin of cosmological entropy. This problem was resolved \\cite{2} by taking into account the cosmological consequence of irreversible particle creation phenomena in the framework of Einstein's equation, classically.\\\\  \\noindent Particle creation phenomena was explored largely during the last century to explain the early Universe. Matter constituents may be produced quantum mechanically \\cite{p}, \\cite{bd}, \\cite{mw} in the framework of Einstein's equations. Cosmological consequence of particle creation mechanism is studied taking into account an explicit phenomenological balance law (see appendix) for the particle number \\cite{2}, \\cite{3}, \\cite{l} in addition to the familiar Einstein's equations. In view of such a balance law, Prigogine et-al \\cite{2} successfully explained the cosmological evolution of the early Universe.  \\\\  \\noindent Recently, Lima et-al \\cite{l:s} have developed a late time model Universe, taking into account the creation phenomena in the matter dominated era. The model admits early deceleration followed by a recent acceleration of the Universe as suggested by present observations and fits SNIa data to some extent. However, they \\cite{l:s} have shown that in their model the creation phenomena is never ending, as a result cosmic evolution does not ever include standard radiation dominated or matter dominated Friedmann era. This definitely creates problem in explaining the structure formation of the Universe and the CMBR. Later, Steigman et-al \\cite{g:r} analyzed the model in the two limits of high and low redshifts. They have observed a clear conflict between the WMAP constraint on matter-radiation equality $z_{eq}$ at high redshift and SNIa data at low redshift. The main criticisms of the $\\beta-\\gamma$ model proposed by Lima et-al \\cite{l:s} are that they have not taken into account the amount of CDM created in the very early Universe at one end, and that their creation rate $\\Gamma = 3\\beta H + 3\\gamma H_0 $ depends on the present Hubble parameter, on the other. If these problems are alleviated, then phenomenologically particle creation process obviously unifies early inflation with late stage of cosmic acceleration in an elegant fashion. \\\\   \\noindent In the present work, we propose a model where, instead of choosing the creation parameter $\\Gamma$ arbitrarily, we have considered the experimentally verified fact that the Universe has recently entered an accelerated phase of expansion. As a result we have chosen the scale factor judiciously, such that particle creation could start again in the matter dominated era. This naturally alleviates the said problems and unifies the early and the late stages of cosmic evolution, in an elegant fashion. Additionally, the model fits perfectly with the WMAP constraint on matter-radiation equality $z_{eq}$ only if one considers the presence of nearly $26\\%$ of primeval matter in the form of baryons and CDM. In view of such a model, particle creation phenomena is now able to explain the history of cosmic evolution from the very early Universe till date, without requiring dark energy at any stage and thus avoiding coincidence problem.\\\\   \\noindent In the following section we write down the field equations incorporating the phenomena of particle creation, which is apparent through the balance equation. In section 3, we show how the conflict between high redshift and the low redshift data may be reduced in the model presented by Lima et al \\cite{l:s}, just by accounting for some amount of CDM produced in the very early Universe. However, we also mention some more problems which are still associated with their \\cite{l:s} model. In section 4, instead of choosing a form of the creation rate arbitrarily, we rather choose a form of the scale factor, suitable for a transition from the early deceleration to the late time cosmic acceleration, to understand the associated problems. This gives us insight to find a form of the creation rate $\\Gamma$ associated with the so called intermediate inflation \\cite{if}. Such a form of $\\Gamma$ appears to be much elegant to study the cosmological evolution alleviating all the problems discussed. This has been done in section 5. Finally we end up with the conclusion in section 6 and an appendix in section (7), to calculate the balance law.\\\\  \\section {Balance law and the Field equations}  As mentioned in the introduction, cosmological consequence taking into account particle creation phenomena is studied by using an explicit phenomenological balance law for the particle number \\cite{2},\\cite{3}. Such a balance law in the process of particle production is modeled by $(nu^{\\alpha})_{;\\alpha}=\\Psi,$ with a source term $\\Psi$. For a vanishing $\\Psi$ the particle number is conserved. In the isotropic and homogeneous background metric,  \\be ds^2=-dt^2+a(t)^2\\left[\\frac{dr^2}{1-kr^2}+r^2 (d\\theta^2+sin^2\\theta d\\phi^2)\\right] \\ee  \\noindent above balance law reduces to, $\\frac{\\dot N}{N} = \\frac{\\dot n}{n}+\\Theta = \\frac{\\Psi}{n} = \\Gamma,$ where, $u^{\\alpha}$ , $N$ , $n$ and $H$ are the fluid four velocity vector, the total number of particles, particle number density and the Hubble parameter respectively, while, $\\Theta = u^{\\alpha}\\;_{;\\alpha} = 3 H$ is the expansion scalar. Thus, the field equations in the spatially flat ($k = 0$) Robertson Walker metric (1) associated with particle creation phenomena are,  \\be 2\\dot H + 3H^2 = - 8\\pi G(p_m+p_{cm}),\\ee  \\be 3H^2  =  8\\pi G (\\rho_m + \\rho_{cm}) = 8\\pi G \\rho,\\ee  \\be 3H + \\frac{\\dot n}{n} = \\Gamma = \\frac{\\Psi}{n},\\ee  \\be p_{cm} = -\\frac{\\rho + p_m}{3H}\\Gamma.\\ee \\noindent For the last equation (5) please see the appendix. In the above set of equations, $H = \\frac{\\dot a}{a}$ is the Hubble parameter, $p_m$ and $\\rho_m$ are the pressure and the energy density of the matter existing in the Universe in the form of a barotropic fluid containing baryons and cold dark matter, created in the very early Universe. $p_{cm}$ and $\\rho_{cm}$ are the pressure and the energy density of the cold dark matter in the form of WIMP created at the late stage of cosmic evolution, ie., during matter dominated era. $\\Gamma$ is the creation parameter, and $n$ is the particle number density. It has been shown in the appendix that the second law of Thermodynamics allows the creation of particle from the gravitational field and the process is irreversible. Thus the creation parameter $\\Gamma > 0$, and so it is clear from equation (5) that the particle creation phenomena is always associated with a negative pressure $p_{cm}$, which may be responsible for acceleration at the late stage of cosmic evolution. We would like to mention at this stage that, other than cold dark matter and baryons, different types of particles may be created in view of quantum field theory in curved space time \\cite{p}, \\cite{bd}, \\cite{mw} and all are associated with a negative pressure, as mentioned. However, also, as mentioned in the introduction that the structure formation requires $22\\%$ that are constituted by dark matter, which must have been created in the early Universe. The present estimated amount of baryons is $4\\%$ and the rest amount required for present cosmic acceleration is usually treated as dark energy. Other form of matter (like hot dark matter, say) has negligible contribution. Here, we proceed to show that creation of the same amount of CDM also solves the puzzle. Further, note that the creation of baryons and CDM in the early Universe was also associated with a large negative pressure. However, the creation phenomena weakened and finally stopped, when Universe expanded sufficiently, thereby giving way to the hot big bang followed by the radiation dominated era of Friedmann type ($a \\propto t^{\\frac{1}{2}}$) \\cite{2}. There after, the baryons and the CDM created in the very early Universe obviously start acting as pressureless dust, so that we can take $p_m = 0$ and hence $\\rho_m = \\rho_0 a^{-3}$, $\\rho_0$ being a constant. Thus we can simplify the above equations to get,  \\be \\Gamma = 3 H + 2\\frac{\\dot H}{H} = 3H\\left(\\frac{2\\dot H+3H^2}{3H^2}\\right) = -3Hw_{e},\\ee  \\be \\rho_{cm} = \\frac{3H^2}{8\\pi G} - \\rho_0 a^{-3},\\ee  \\be p_{cm} = - \\frac{1}{8\\pi G}H\\Gamma,\\ee  \\noindent $w_e$ being the effective state parameter. In view of the above set of three equations (6) - (8), we need to find the scale factor $(a)$ (and consequently $(H)$, the Hubble parameter), the creation rate $\\Gamma$, the creation pressure $p_{cm}$ and the creation matter density $\\rho_{cm}$. Obviously, we need yet another suitable condition to solve the system of equations. Lima et al \\cite{l:s} studied these equations under the assumption of a form of the creation rate $\\Gamma$. Afterwards, while further studying their model in connection with data fitting, they found a clear conflict between the low (SNIa) and high (WMAP constraint on $z_{eq}$) \\cite{g:r} redshift limits. In the following section, we review the problem and show how the production of CDM in the very early Universe alleviates the problem. ",
        "conclusions": "} We have presented a phenomenological cosmological model based on particle creation in the matter-dominated era, which  fits SNIa data and the redshift of the early integrated Sachs-Wolfe effect at the matter-radiation equality, $z_{eq} = 3145^{+140}_{-139}$ \\cite{w1} and $z_{eq} = 3196^{+134}_{-133}$ \\cite{w2}, determined by WMAP. The value of h-parameter $(h\\approx 0.70)$ and $H_0 t_0 \\approx 1$ also are very much within observational limit obtained from HST project \\cite{hst}. Such a model constraints the amount of primeval matter to $0.24\\le\\Omega_m\\le0.26$, out of which CDM amounts to $0.20 \\le \\Omega_{CDM}\\le 0.22$, which is again the same amount required for structure formation.\\\\  \\noindent Quantum particle production phenomena (QFT in CST) has been discussed in the literature in detail \\cite{bd}, \\cite{mw}. The energy of these particles may then be extracted from the gravitational field \\cite{1}. To study the classical consequence of particle creation phenomena, kinetic collision theory may be adopted to find a balance law in addition to standard Einstein's equations, if weakly interacting massive particles (WIMP) are taken into account. In view of such a balance law a cosmological model of the early Universe has been explored \\cite{2}. In that model, as the particle production rate becomes comparable to the expansion scalar ($\\Theta = 3H$), it builds up a large negative creation pressure that pushes the Hubble parameter $H$ to approximately constant value. As a result, inflationary behaviour due to a large particle production rate is realized. Consequently, the universe starts with a de Sitter phase avoiding cosmological singularity. As inflation continues, the expansion rate becomes too large comparable to the particle production rate. In such a dilute cosmic fluid, particle production rate decreases and halts at some time. At this stage, in the absence of sufficient negative pressure, inflation ends giving way to reheating and the resulting Universe smoothly approaches the familiar Friedmann-Lemaitre-Robertson-Walker behaviour. Since the radiation dominated era, in such a model is of standard Friedmann type (i.e., $a \\propto t^{1/2}$, $a$ being the scale factor), so the standard Big-Bang-Nucleosynthesis (BBN) remains unaltered. Thus cosmological evolution of the early Universe may be explained successfully in view of particle creation phenomena.\\\\  \\noindent Inflation makes the Universe almost spatially flat, which has also been confirmed by recent observations. However, at the end of inflation, if the Universe remains slightly away from spatial flatness, it may cause particle production in the matter dominated era again. If this happens, then even a slow particle production rate may cause sufficient negative pressure in billions of years to cause recently observed cosmic acceleration. This phenomena has been studied earlier \\cite{j}, \\cite{l:s}, but the model suffered from a clear conflict \\cite{g:r} between the low (SNIa) and high redshift (WMAP) data. This problem has been resolved in the present model, which constraints primeval matter to $\\Omega_m \\approx 26\\%$. Since, particle production starts long after matter-radiation equality, so the early growth of perturbation in connection with the structure formation also remains unaltered following the $\\Lambda$CDM concordance model closely. Thus particle production process may successfully explain the late stage of cosmic evolution also. As a result, the cosmological evolution of the Universe may be explained successfully in view of particle creation phenomena from early time till date.\\\\  \\noindent Presently we are having numerous dark energy models, explaining the late time cosmic phenomena. Most of these models do not explain the early Universe at one hand and it is not possible to identify these models from one another in any of the future experiments, on the other. Cosmological consequence of particle creation phenomena does not require dark energy at all at any stage, and some programmes have been taken in the recent years to detect lightest neutralino of roughly $10-10000$ GeV - one of the weakly interacting massive particles (WIMP). Thus, if creation phenomena of cold-dark-matter can solve the presently observed cosmic puzzle single handedly, without taking into account the dark energy at all, may be resolved in near future.\\\\"
    },
    "1005/1005.4208_arXiv.txt": {
        "abstract": "Extended \\hi\\ structures around galaxies are of prime importance to probe galaxy formation scenarios. The giant \\hi\\ ring in the Leo group is one of the largest and most intriguing \\hi\\ structures in the nearby universe. Whether it consists of primordial gas, as suggested by the apparent absence of any optical counterpart and the absence of an obvious physical connection to nearby galaxies, or of gas expelled from a galaxy in a collision is actively debated.  We present deep wide field-of-view optical images of the ring region obtained with MegaCam on the CFHT. They reveal optical counterparts to several \\hi\\ and UV condensations along the ring, in the $g'$, $r'$, and $i'$ bands, which likely correspond to stellar associations formed within the gaseous ring.  Analyzing the spectral energy distribution of one of these star-forming regions, we found it to be typical for a star-forming region in pre-enriched tidal debris. We then use simulations to test the hypothesis that the Leo ring results from a head-on collision between Leo group members NGC~3384 and M96. According to our model which is able to explain, at least qualitatively, the main observational properties of the system, the Leo ring is consistent with being a collisional ring. It is thus likely another example of extended intergalactic gas made-up of pre-enriched collisional debris. ",
        "introduction": "The quest for primordial gas clouds that have never been involved in star-forming episodes in the local universe has motivated extensive \\hi\\ surveys (e.g., HIPASS, \\citealp{hipass}; and ALFALFA, \\citealp{alfalfa}) but has yet been rather unsuccessful. Among the putative candidates are the so-called dark galaxies \\citep[e.g.,][]{Davies04}, which are however likely tidal debris from galaxy interactions \\citep{Duc08}. Isolated \\hi\\ clouds around early-type galaxies \\citep[ETGs;][]{Morganti06,Oosterloo07,Serra10} could either be accreting clouds or the collisional remnants of the violent merger events that are usually invoked for the formation of ETGs. Investigating the stellar populations associated or not with these clouds could discriminate these formation scenarios.  The Leo ring, in the Leo group of galaxies, is one of the most spectacular and mysterious intergalactic \\hi\\ structures known in the nearby universe \\citep{Schneider83,Schneider85,Schneider89b,Stierwalt09}.  It has a ring-like shape of diameter $\\sim 200$~kpc, quite asymmetric and somewhat clumpy, apparently centered on the NGC~3384/M105 galaxy pair. Some radial filamentary structures are observed, in particular a bridge connecting the ring to the spiral galaxy M96.  Numerous attempts were done to find counterparts to the \\hi\\ emission at other wavelengths, which remained unsuccessful in the optical and H$\\alpha$ regime. This had suggested that the ring could consist of primordial gas. The only relatively bright optical counterparts correspond to \\hi-rich dwarf galaxies in the group \\citep{Stierwalt09} whose velocities are discrepant from the ring velocity field (e.g., HGC 201970)\\footnote{One exception is HGC 202027, whose Sloan Digital Sky Survey (SDSS) redshift is close to that of the \\hi\\ ring, but is not associated with an \\hi\\ clump peak of the ring. It is most likely a pre-existing group member unrelated to the formation of the \\hi\\ ring.}.  Lastly, \\citet{Thilker09} reported a UV counterpart to some \\hi\\ clumps along the ring from \\textit{Galaxy Evolution Explorer (GALEX)} observations. The UV knots with the highest far-UV to near-UV flux ratios are good candidates for being star-forming regions associated with the ring (and not background galaxies). The inferred star formation rate is very low but consistent with the \\hi\\ column densities below $2 \\times 10^{20}$~cm$^{-2}$. Based on the FUV--NUV color and upper limits on the FUV--r color, \\citet{Thilker09} estimated a very low gas metallicity, and concluded that the \\hi\\ ring is likely primordial.  \\begin{figure*}[th] \\begin{center} \\includegraphics[width=0.9\\textwidth]{figure1.eps} \\caption{Leo ring. Left: distribution of the \\hi\\ gas, as mapped at Arecibo (reddish contours; from \\citealt{Schneider89}), the WSRT (in brown; T. A. Oosterloo et al., in preparation), superimposed on a deep optical image obtained at the CFHT.  Right: zoom on several UV condensations \\citep{Thilker09}. The \\textit{GALEX} FUV (white contours) and WRST \\hi\\ emission (brown contours) are superimposed on the optical true color image made from $g'$, $r'$, and $i'$ bands data. The apertures used by \\cite{Thilker09} are shown in green. } \\end{center} \\label{fig:megacam} \\end{figure*}  More recently, \\cite{Bot09} reported a marginal detection of dust emission at 8 and 24 $\\mu$m toward the densest \\hi\\ condensation. If real, it implies that the metallicity of the \\hi\\ ring is higher than the $Z_{\\odot}$/20 estimated by \\cite{Thilker09}. This would support scenarios where the gas has a galactic origin. Indeed, clouds expelled during galaxy interaction have been previously processed in the parent galaxies and are thus relatively metal-rich. \\cite{Bekki05b} proposed that the disruption of a low surface brightness galaxy in the gravitational potential of a group forms gaseous arcs and rings with column density in agreement with observations, without providing a specific model for the Leo ring, though.  Collisional rings are a more classical example of such debris \\citep{Appleton96,Bekki98,Bournaud03a}: a head-on galaxy collision disrupts a gaseous galactic disk into an expanding pre-enriched ring of gas, whose star-forming activity depends on the remaining gas density.  In this Letter, we present new optical images of the system and report the detection of counterparts in the $g'$, $r'$, and $i'$ bands (Section 2). We present a numerical model of the Leo ring as the result of a head-on collision in the Leo group (Section 3) and discuss this hypothesis in Section 4. ",
        "conclusions": "\\label{sec:disc}  The origin of the Leo ring has been actively debated for more than two decades. If made of primordial gas, it would be the only such giant \\hi\\ structure known in the nearby universe. The absence of an optical counterpart was often suggested to preclude a collisional origin. Our deep optical images have confirmed that indeed it is not associated with a diffuse stellar component brighter than 28 mag~arcsec$^{-2}$. However, as shown by our numerical model of the Leo ring and previous simulations \\citep{Bournaud07}, collisional rings are not expected to systematically contain a significant populations of old stars expelled from the progenitor galaxy. Old stars can spread in a halo rather than along the ring, as is the case in our model with a predicted brightness fainter than $\\sim$ 29~mag~arcsec$^{-2}$.   The Leo ring is striking for not having managed to convert a significant fraction of its gas into stars. If primordial, this would require unexampled long-term stability conditions \\citep[stabilization by a deep potential such as the Leo group could nevertheless quench star formation;][]{martig09}. Our imaging program has revealed the presence of compact, optical counterparts to the far-ultraviolet sources detected by \\textit{GALEX}. These are likely young stellar associations formed in situ in the \\hi\\ ring, with very low star formation rates that are consistent with the low \\hi\\ surface densities. \\cite{Thilker09} had estimated a very low metallicity from UV data, suggesting the presence of primordial gas. Adding three data points in the optical bands to the SED of one of the UV clumps, we find it most consistent with that of a star-forming object in the NGC~5291 ring, known to be pre-enriched and of collisional origin \\citep{Bournaud07}. Thus, available data are actually consistent with the Leo ring being made-up of pre-enriched material.  Obtaining a direct measurement of its metallicity would give the final word.    If the Leo ring is not made-up of primordial gas, its formation should involve pre-enriched gas expelled during galaxy--galaxy or galaxy--group interactions. Ram pressure would require a very dense intracluster medium to produce such a massive and extended structure, a condition not met in a relatively loose group like Leo. Note also that ram pressure usually forms one-sided tails attached to the parent galaxy, not ring-like structures \\citep{tonnessen09}. The tidal potential of the group could also affect the gas content of galaxies, but is generally a second-order effect compared to direct collisions in galaxy pairs.   A more likely hypothesis is that the Leo ring has the same origin as the confirmed collisional ring around NGC~5291 \\citep{Longmore79,Bournaud07} or the famous Cartwheel ring \\citep{Horellou01} --- the later being much denser and actively star forming.  Our numerical model reproduces the main characteristics of the Leo ring, after a galaxy collision at a relatively moderate speed, quite plausible in the Leo group. The proposed progenitor galaxy, NGC 3384, does not show strong morphological disturbances, even in our deep imaging data. This is however expected in our model, that produces a gaseous collisional ring, an extended and very faint (unobservable) stellar halo, but leaves little signs of collision such as classical tails and shells around the parent galaxy. \\texttt{SAURON} observations of NGC~3384 \\citep{Emsellem04,Kuntschner06} revealed a barred rotating disk, with a low gas and dust content and without very recent star formation, which is also broadly consistent with our model where the pre-existing spiral is converted into a fastly rotating ETG by the 1.2 Gyr old collision.  Large collisional rings can form in groups of galaxies and take an appearance a priori suggesting primordial gas rather than tidal debris. The Leo ring appears fully consistent with being the result of such a process."
    },
    "1005/1005.0271_arXiv.txt": {
        "abstract": "We present the angular correlation function of the X-ray population of 1063 XMM-Newton observations at high Galactic latitudes, comprising up to $\\sim$30000 sources over a sky area of $\\sim$125 sq. degrees in the energy bands: soft (0.5-2 keV) and hard (2-10 keV). This is the largest sample of serendipitous X-ray sources ever used for clustering analysis purposes to date and the results have been determined with unprecedented accuracy. We detect significant clustering signals in the soft and hard bands ($\\sim$10$\\sigma$ and $\\sim$5$\\sigma$, respectively). We deproject the angular correlation function via Limber's equation and calculate the typical spatial lengths. We infer that AGN at redshifts $\\sim$1 are embedded in dark matter halos with typical masses of $\\log M \\sim 12.6 h^{-1}$~$M_{\\odot}$ and lifetimes in the range $\\sim3-5 \\times 10^8$~years, which indicates that AGN activity is a transient phase in the life of galaxies. ",
        "introduction": "Active galactic nuclei (AGN) are the brightest persistent extragalactic sources known, with their X-ray emission the most common feature among them. Thanks to their large bolometric output, AGN can be detected through cosmological distances, which makes them essential tracers of galaxy formation and evolution, as well as the large-scale structure of the Universe. Clustering studies of AGN at redshift $\\sim$1, when strong structure formation processes took place, are key tools for understanding the underlying mass distribution and evolution of cosmic structures. ",
        "conclusions": "\\label{conclusions}  We have studied the angular correlation function of a large sample of serendipitous X-ray sources from 1063 XMM-Newton observations at high Galactic latitudes in the 0.5-2~keV (soft), and 2-10~keV (hard) energy bands, respectively. Our sample comprises 31288 and 9188 sources in the soft and hard bands, respectively, covering $\\sim$125.5~deg$^{2}$ in the sky, thus being the largest sample ever used in clustering investigations.  We found significant positive angular clustering signal in the soft ($\\sim$10$\\sigma$) and hard ($\\sim$5$\\sigma$) bands. The result in the hard band clears up the debate on whether X-ray sources detected in this band cluster or not, since a number of past works had reported different inconclusive results ranging from a few $\\sigma$ detections to no detection at all.  We inverted Limber's equation, assuming a given redshift distribution for our sources, in order to estimate typical spatial correlation lengths. We used these values to calculate the $rms$ fluctuations of the AGN distributions within a sphere of radius 8 $h^{-1}$~Mpc, and compared them with that of the underlying mass distribution from the linear theory in order to estimate the bias parameter of our X-ray sources. The bias depends on the mass of the dark matter haloes (DMH) that host the AGN population. From the computed bias values we have estimated a typical DMH mass of $\\langle \\log M_{DMH}\\rangle \\simeq 12.60\\pm0.34$~$h^{-1}$~M$_\\odot$.  The typical AGN lifetime derived from the Press-Schechter approximation at redshift $z \\sim 1$ lies in the range in the range $t_{AGN}=3.1-4.5 \\times 10^8$~yr. This interval is significantly shorter than the time span between that redshift and the present thus suggesting the existence of many AGN generations, and that a significant fraction of galaxies may switch from a quiescent phase to AGN activity, and vice versa, several times throughout their lifes."
    },
    "1005/1005.0876_arXiv.txt": {
        "abstract": "The last decade has shown us that the observational properties of neutron stars are remarkably diverse.  From magnetars to rotating radio transients, from radio pulsars to `isolated neutron stars,' from central compact objects to millisecond pulsars, observational manifestations of neutron stars are surprisingly varied, with most properties totally unpredicted.  The challenge is to establish an overarching physical theory of neutron stars and their birth properties that can explain this great diversity.  Here I survey the disparate neutron stars classes, describe their properties, and highlight results made possible by the {\\it Chandra X-ray Observatory}, in celebration of its tenth anniversary. Finally, I describe the current status of efforts at physical `grand unification' of this wealth of observational phenomena, and comment on possibilities for {\\it Chandra}'s next decade in this field. ",
        "introduction": "The {\\it Chandra} era has seen the proliferation of a greater variety of possibly distinct observational classes of neutron star than ever before (not even including accreting sources, ignored in this review). With emission spanning the electromagnetic spectrum and radiative properties that span a huge fraction of conceivable phase space, such incredible range and diversity is not only unpredicted, but in many ways astonishing given the perhaps naively simple nature of the neutron star.  Collapsed and close cousins to black holes, why should neutron stars exhibit so much `hair'?  Indeed, the sheer number of different class names is confusing even within the field.  We therefore begin by presenting an introductory census of the currently identified neutron zoo.  The classic young {\\bf `radio pulsars'} (PSRs) pulse regularly and predictably across the EM spectrum though generally are most observable in the radio band.  The PSR prototype is often identified as the Crab pulsar, even though it is far from prototypical of the class, given its unparalleled energy output ($10^{38}$ erg/s), and its magnificent nebula (see Figure 1).  The `radio pulsar' term is not observationally accurate for the class, as, for example the `Geminga' pulsar has all the properties of a radio pulsar except observable radio emission, possibly but not certainly due to unfortunate beaming geometry \\cite{hh92}.  The term {\\bf `rotation-powered pulsar'} (RPP) is therefore more cautious and precise: these objects are powered by their loss of rotational energy due to braking by their magnetic fields.  {\\bf `Millisecond pulsars'} (MSPs), though also rotation-powered, have different evolutionary histories, involving long-lived binary systems and a `recycling' accretion episode which spun-up the neutron star and quenched its magnetic field \\cite{bv91}.  In this review, MSPs are a sub-class of RPPs and have similar emission properties, although {\\it Fermi} has recently shown that MSPs are surprisingly bright $\\gamma$-ray sources, suggesting refinements to high-energy emission models are required \\cite{aaa+09}.  The very recently discovered {\\bf `rotating radio transients'} (RRATs) \\cite{mll+06} do not seem to produce observable periodic emission, the defining property of RPPs.  Rather, they produce unpredictable sudden short radio bursts, which occur at integral multiples of an underlying periodicity. These may be a sub-class of RPPs; this is discussed further below.  Radio emission was long thought to be the hallmark of the non-accreting neutron star, Geminga notwithstanding. But several different sub-classes of generally radio-quiet neutron star have emerged in the {\\it Chandra} era.  The {\\bf `isolated neutron stars'} (INS; poorly named as most RPPs are also isolated but are not `INSs') have as defining properties quasi-thermal X-ray emission with relatively low X-ray luminosity, great proximity, lack of radio counterpart, and relatively long periodicities ($P=$3--11~s).  Then there are the ``drama queens'' of the neutron-star population:  the {\\bf `magnetars'}.  Magnetars have as their true defining properties occasional huge outbursts of X-rays and soft-gamma rays, as well as luminosities in quiescence that are generally orders of magnitude greater than their spin-down luminosities.  Magnetars are thought to be young, isolated neutron stars powered ultimately by the decay of a very large magnetic field. At their ``baddest,'' magnetars can briefly outshine all other cosmic soft-gamma-ray sources {\\it combined} \\cite{hbs+05}.  The above census would be incomplete without mention of the handful of X-ray bright {\\bf `compact central objects'} (CCOs), a likely heterogeneous class. They are so-named because of their central location in supernova remnants (SNRs), and have otherwise, until very recently, baffling properties.  The challenge of the past decade was -- and continues to be -- to find a way to unify this variety into a coherent physical picture. What determines whether a neutron star will be born with, for example, magnetar-like properties or as a Crab-like pulsar?  What are the branching ratios for the various varieties, and, given estimates of their lifetimes, how many of each are there in the Galaxy? Ultimately such questions are fundamental to understanding the fate of massive stars, and the nature of core-collapse, while simultaneously relating to a wider variety of interesting high-energy astrophysics, ranging from the equation-of-state of ultradense matter, to the physics of matter in ultrahigh magnetic fields. In this contribution, we review the above observationally defined classes, and consider the status of the progress toward `grand unification' of neutron stars into a coherent physical understanding of the births and evolution of these objects. Emerging as central to the picture are the stellar surface dipolar magnetic fields (estimated via $B=3.2 \\times 10^{19} (P \\dot{P})^{1/2}$~G, where $\\dot{P} = dP/dt$) and ages (estimated via the `characteristic age' $\\tau \\equiv P/2\\dot{P}$). We show how observations at X-ray energies are essential to this work, and highlight, as much as space permits, some of the great contributions made by {\\it Chandra} this past decade. ",
        "conclusions": ""
    },
    "1005/1005.4791_arXiv.txt": {
        "abstract": "{} {NGC~2992 is an intermediate Seyfert 1 galaxy showing outflows on kilo parsec scales which might be due either to AGN or starburst activity. We therefore aim at investigating its central region for a putative starburst in the past and its connection to the AGN and the outflows.} {Observations were performed with the adaptive optics near infrared integral field spectrograph SINFONI on the VLT, complemented by longslit observations with ISAAC on the VLT, as well as N- and Q-band data from the Spitzer archive. The spatial and spectral resolutions of the SINFONI data are 50\\,pc and 83\\,km\\,s$^{-1}$, respectively. The field of view of $3\\arcsec\\times 3\\arcsec$ corresponds to 450\\,pc $\\times$ 450\\,pc. Br$\\gamma$ equivalent width and line fluxes from PAHs were compared to stellar population models to constrain the age of the putative recent star formation. A simple geometric model of two mutually inclined disks and an additional cone to describe an outflow was developed to explain the observed complex velocity field in H$_2$\\,1-0S(1).} {The morphologies of the Br$\\gamma$ and the stellar continuum are different suggesting that at least part of the Br$\\gamma$ emission comes from the AGN. This is confirmed by PAH emission lines at 6.2\\,$\\mu$m and 11.2\\,$\\mu$m and the strength of the silicon absorption feature at 9.7\\,$\\mu$m, which point to dominant AGN activity with a relatively minor starburst contribution. We find a starburst age of 40\\,Myr -- 50\\,Myr from Br$\\gamma$ line diagnostics and the radio continuum; ongoing star formation can be excluded. Both the energetics and the timescales indicate that the outflows are driven by the AGN rather than the starburst. The complex velocity field observed in H$_2$\\,1-0S(1) in the central 450\\,pc can be explained by the superposition of the galaxy rotation and an outflow. } {} ",
        "introduction": "\\object{NGC 2992}, initially classified as a Seyfert 2, was later classified as an intermediate Seyfert 1 galaxy on the basis of a broad H$\\alpha$ component with no corresponding H$\\beta$ component in its nuclear spectrum (Ward et al. \\cite{ward80}). Later, Glass (\\cite{glass97}) suggested that it is a hybrid between an intermediate Seyfert and a starburst galaxy, induced by the interaction with NGC 2993 to the south. Both galaxies are connected by a tidal tail with a projected length of 2\\farcm9.  NGC 2992 is highly inclined by about 70$\\degr$ to our line of sight with a broad disturbed lane of dust in its equatorial plane. The redshift is 0.0077 which corresponds to a distance of 32.5\\,Mpc and an angular scale of 150\\,pc per arcsec assuming $H_0$=75\\,km\\,s$^{-1}$\\,Mpc$^{-1}$.  At a wavelength of 6\\,cm NGC 2992 shows prominent loops to the northwest and southeast. This feature, designated as figure-of-8, has a size of about 8\\arcsec\\ with the long axis at a position angle (PA) of $-26\\degr$, which is not coincident or perpendicular to the major axis (Ulvestad \\& Wilson \\cite{ulvestad84}). Chapman et al. (\\cite{chapman00}) conclude from their IR observations that the most convincing interpretation for these loops is that of expanding gas bubbles which are seen as limb-brigthened loops as was already proposed by Wehrle \\& Morris (\\cite{wehrle88}). Two possible explanations are discussed for the origin of the loops: They may be driven by an AGN or a superwind from a compact starburst. X-ray observations point to the first possibility: Colbert et al. (\\cite{colbert98}) concluded that the large scale outflows are driven by nonthermal jets from the AGN that entrain the surrounding material and heat it by shocks over kpc spatial scales.  Heckman et al. (\\cite{heckman81}) reported complex kinematics in NGC 2992 on the basis of longslit spectra at PA=120\\degr. This work was extended by Colina et al. (\\cite{colina87}) who find indications for an outflow within a plane highly inclined to the plane of the galaxy. M\\'arquez et al. (\\cite{marquez98}) measured [\\ion{O}{iii}] and H$\\alpha$ lines at nine different position angles and concluded, although from a simple kinematical model, that the double peaked line profiles and asymmetries can be accounted for by an outflow within a conical envelope or on the surface of a hollow cone. The high spectral resolution observations (4.9\\,\\AA: 3440\\,\\AA\\ -- 5579\\,\\AA, 11.9\\,\\AA: 5840\\,\\AA\\ -- 9725\\,\\AA) of narrow line regions in NGC 2992 of Allen et al. (\\cite{allen99}) show at most positions a double component line profile. One component of the double peaked lines follows the galactic rotation curve, while the other is identified as a wind that is expanding out of the plane of the galaxy with a velocity of up to 200\\,km\\,s$^{-1}$. These authors did not find spectral evidence for a significant contribution from a starburst to the outflow.  Depending on the inclinations of galaxy and cone several interpretations of these kinematics are possible. It is generally assumed that the NW edge of the disk of NGC 2992 is closest to us with trailing spiral arms, and that the SE cone is closer and directed at us. However, an arc-shaped emission found by Garc\\'ia-Lorenzo et al. (\\cite{garcia01}), which might be associated with the radio loop, is red shifted in the SE, and can therefore only be consistent with an outflow if it is pointing away from us.  In this paper, we present spatially and spectrally resolved K-band observations of the inner $3\\arcsec\\times3\\arcsec$ of NGC 2992 taken with the adaptive optics integral field spectrograph SINFONI. These are complemented by longslit observations along the major axis out to a radius of about 5\\arcsec\\ with ISAAC, and N- and Q-band spectra from the Spitzer archive. SINFONI data were taken with unprecedent spatial resolution in the K-band which allows us to study the dynamics of gas and stars in the central region in detail. Together with the spectral resolution of SINFONI this enables us to search for a starburst and its connection to the AGN and, possibly, to the outflows. After presenting the observational results in section 3, we discuss the observations in section 4, and show that there is evidence for a nuclear starburst in NGC 2992. Finally, in section 5 we concentrate on the dynamics of the gas and stars and the consequences on the geometry of the galaxy. ",
        "conclusions": "We have presented a detailed analysis of near infrared K-band spectra complemented by N- and Q-band spectra of NGC 2992. K-band spectra were obtained with the adaptive near infrared integral field spectrograph SINFONI and allow a reconstruction of the 2-dimensional distribution and kinematic of the stars and gas in the inner $3\\arcsec\\times 3\\arcsec$ (450\\,pc) at an angular resolution of 0\\farcs3. The N- and Q-band data were obtained with Spitzer and taken from the archive. We compared the equivalent width of Br$\\gamma$ and the supernova rate, which was derived from radio data from the literature, to STARS evolutionary synthesis models and find evidence for a short burst of star formation $(40 - 50)$\\,Myr ago.  From our near-infrared data as well as equivalent widths and luminosities of fine-structure emission lines and polycyclic aromatic hydrocarbon features detected in the N- and Q-band spectra we were able to estimate the nuclear star formation rate and quantify the contribution of the star formation to the AGN luminosity: the luminosity is dominated by the AGN activity with a contribution of only $10\\%-20\\%$ from nuclear star formation. This is in agreement with the conclusion that part of the Br$\\gamma$ emission cannot be attributed to star formation but is due to the AGN.  Energy and timescale considerations let us conclude, that the starburst would have to be more recent and more intense to drive the figure-of-8 as well as the large scale (2\\,kpc) outflow. On the other hand the estimation of the luminosity of the AGN of $10^{44}$\\,erg\\,s$^{-1}$ is about an order of magnitude higher than is necessary to drive the outflows.  The observed velocity curve in H$_2$\\,1-0S(1) of NGC 2992 can be described as the superposition of the galaxy rotation and a conical outflow with constant velocity. The different position angles of the inner and outer region can be modelled by disks with different orientations in space and thus being probably dynamically decoupled. However, these disks are not related to the starburst. In this picture the south east cone is pointing away from us supporting the findings from Garc\\'ia-Lorenzo (\\cite{garcia01})."
    },
    "1005/1005.1661_arXiv.txt": {
        "abstract": "We present UV luminosity functions (LFs) at 1500 \\AA\\ derived from the HST Early Release Science WFC3/UVIS data acquired over $\\sim50$ arcmin$^2$ of the GOODS-South field. The LFs are determined over the entire redshift range $z=0.75-2.5$ using two methods, similar to those used at higher redshifts for Lyman Break Galaxies (LBGs): (1) 13-band UV+optical+NIR photometric redshifts to study galaxies in the range $z=0.5-2$ in three bins of $dz=0.5$, and (2) dropout samples in three redshift windows centered at $z\\sim1.5$, $z\\sim 1.9$, and $z\\sim 2.5$.  The characteristic luminosity dims by 1.5 mag from $z=2.5$ to $z=0.75$, consistent with earlier work.  However, the other Schechter function parameters, the faint-end slope and the number density, are found to be remarkably constant over the range $z=0.75-2.5$. Using these LF determinations we find the UV luminosity density to increase by $\\sim1.4$ dex according to $(1+z)^{2.58\\pm0.15}$ from $z\\sim0$ to its peak at $z\\sim2.5$.  Strikingly, the inferred faint-end slopes for our LFs are all steeper than $\\alpha=-1.5$, in agreement with higher-redshift LBG studies. Since the faint-end slope in the local universe is found to be much flatter with $\\alpha\\simeq-1.2$, this poses the question as to when and how the expected flattening occurs. Despite relatively large uncertainties, our data suggest $\\alpha\\simeq-1.7$ at least down to $z\\sim1$. These new results from such a shallow early dataset demonstrate very clearly the remarkable potential of WFC3/UVIS for the thorough characterization of galaxy evolution over the full redshift range $z\\sim0.5$ to $z\\sim3$. ",
        "introduction": "The UV Luminosity Function (LF) of galaxies is a key diagnostic for establishing the contributions of galaxies of different luminosities and masses to the cosmic star-formation rate (SFR) density. In addition, in combination with other measures, like the galaxy mass function, it also provides essential clues as to the causes for the slow down in star-formation of galaxies in more recent times since $z\\sim2$.  The UV LF has been studied very extensively at redshifts $2\\lesssim z\\lesssim6$, but at lower redshifts ($z\\la2$), similar deep, high resolution data has not been available. To date the $z<2$ UV LF has been measured from GALEX data using the dropout technique \\citep[e.g.][]{Burgarella06,Ly09,Haberzettl09}, and also using spectroscopic or photometric redshifts \\citep[e.g.][]{Arnouts05,Wyder05,Budavari05}.  The main limitations of the GALEX results are (1) the very wide point-spread function (PSF; $\\sim5$\\arcsec\\ FWHM), which may cause blending of the UV light from several sources and makes the identification of optical counterparts relatively difficult, and (2) the limited depth ($\\sim25$ mag).  With the installation of the Wide-Field Camera 3 (WFC3) on the Hubble Space Telescope (HST), the efficient coverage of the electromagnetic spectrum at high spatial resolution ($\\lesssim$0\\farcs15) has been extended to the UV and to the IR.  The new IR capability allowed us (and other groups) to perform the first robust measurements of the UV LF at $z\\sim7-8$ in the first-epoch data of the HUDF09 \\citep[e.g.][]{Oesch10c,Bouwens10a,McLure10,Bunker09}.  The high sensitivity UV imaging capability provided by the WFC3/UVIS channel similarly allows us to derive the UV LF at redshifts below $z\\sim2.5$, where the $\\sim$1500 \\AA\\ UV flux from galaxies could not be probed efficiently with previous high resolution cameras. While the depth of the current ERS WFC3/UVIS UV data \\citep{Windhorst10} is less than that of the WFC3/IR data, the present depths still significantly exceed that from any previous UV data set.  We use two different methods for our LF determinations: (1) 13-band UV+optical+NIR photometric redshifts to study galaxies in the range $z\\sim0.5-2$, and (2) the UV dropout technique, applied to the UVIS filter set, which allows for the selection of star-forming galaxies at $z\\sim1.3$ to $z\\sim2.8$. The latter dropout samples are similar to the recent selections of \\citet{Hathi10}.  In Section \\ref{sec:data} we describe the data, the photometric redshift estimates and the dropout selection. This is followed in Section \\ref{sec:LF} by our constraints on the UV LF between $z = 0.5 - 2.8$, and we end with a short discussion on the evolution of the UV LF parameters and luminosity density in Section \\ref{sec:discussion}.  Throughout this letter, we adopt $\\Omega_M=0.3, \\Omega_\\Lambda=0.7, H_0=70$ kms$^{-1}$Mpc$^{-1}$, i.e. $h=0.7$. Magnitudes are given in the AB system \\citep{Oke83}. ",
        "conclusions": "\\label{sec:discussion}  The evolution of the Schechter parameters of recent UV LFs in the redshift range $z\\sim0-8$ is shown in Figure \\ref{fig:LFparams}.  Our results cover the whole 4.3 Gyr period from $z\\sim0.75$ to $z\\sim2.5$.  Interestingly, in the region of overlap, the LFs we derive from our photometric redshifts and UV dropout selections are in very good agreement. This suggests that our UV dropout samples are well-defined and reasonably complete.  In Figure \\ref{fig:LFparams}, our results are compared to several Schechter function estimates from the literature.  Within the uncertainties, $\\phi_*$ is found to be fairly constant over the entire redshift range, although there is a weak trend towards higher $\\phi_*$ at lower redshifts.   The most noteworthy trend is the monotonic fading of the characteristic luminosity by a factor of $\\sim16$ from $z\\sim3$ ($M_*$ fades from $-21$ at $z\\sim3$ to $-18$ at $z\\sim0$).  % At $z>4$, $M_*$ turns over and decreases monotonically again towards higher redshift, such that $M_*(z=8)$ is essentially equal to $M_*(z\\sim1)$.   It is very interesting that all our faint-end slope measurements are steeper than $\\alpha \\sim -1.5$, even at $z\\sim 0.5-1$, in agreement with the earlier measurements of \\citet{Arnouts05}. Despite relatively large errorbars, the faint-end slope $\\alpha$ is clearly seen to transition from $\\alpha<-1.5$ at $z\\gtrsim0.5-2$ to a flatter $\\alpha=-1.2$ in the local universe \\citep[but see][who find a steeper slope locally]{Treyer98}.  With the present data we cannot determine whether this transition occurs relatively abruptly at low redshift, or smoothly from $z\\sim2$ to $z\\sim0$. %             Finally, in Figure \\ref{fig:LD} we show the cosmic 1500 \\AA\\ luminosity density $\\rho_L$ (uncorrected for dust extinction) estimated from the Schechter function parameters. The luminosity density is shown integrated to $L>0.1L^*_{z=3}$ and also integrated to zero luminosity. Clearly, the evolution of $\\rho_L$ is mainly driven by $M_*$. From the peak at $z\\sim2.5$ the total UV luminosity drops by 1.5 dex to $z\\sim0$. This decrease is well described by $\\rho_L(z)\\propto(1+z)^{\\beta_L}$, for which we find $\\beta_L(z<2.5) = 2.58\\pm0.15$, in good agreement with earlier measurements \\citep[e.g.][]{Schiminovich05,Tresse07}. The evolution at higher redshift follows $\\beta_L(z>2.5) = -1.82\\pm0.37$. $z\\sim2.5$ thus marks the high point in the production of the cosmic UV luminosity and is clearly an important epoch for exploring the properties of galaxies \\citep[e.g. see][for a new technique for selecting galaxies in this range with simple color criteria]{Cameron10}.  The remarkable potential of the HST WFC3/UVIS camera is clear from these new results.  This ERS dataset with its initial rather shallow observations has demonstrated already that WFC3/UVIS is a powerful tool for exploring galaxy evolution over the redshift range $z\\sim0.5$ to $z\\sim3$. % Deeper data will provide key insights into the decline of the cosmic SFRD since $z\\sim2$, particularly at the lowest luminosities. For example, a steep LF with $\\alpha\\sim-1.7$ suggests that star-formation efficiencies and feedback processes are only weakly dependent on halo mass \\citep[e.g.][]{Trenti10}."
    },
    "1005/1005.1511_arXiv.txt": {
        "abstract": "{We present {\\it Herschel}-PACS observations of rest-frame mid-infrared and far-infrared spectral line emissions from two lensed, ultra-luminous infrared galaxies at high redshift: \\object{MIPS J142824.0+352619} (MIPS J1428), a starburst-dominated system at z = 1.3, and \\object{IRAS F10214+4724} (F10214), a source at z = 2.3 hosting both star-formation and a luminous AGN. We have detected [O\\,I]63$\\mu$m and [O\\,III]52$\\mu$m in MIPS J1428, and tentatively [O\\,III]52$\\mu$m in F10214. Together with the recent ZEUS-CSO [C\\,II]158$\\mu$m detection in MIPS J1428 we can for the first time combine [O\\,I], [C\\,II] and far-IR (FIR) continuum measurements for photo-dissociation (PDR) modeling of an ultra-luminous (L$_{IR}\\ge$10$^{12}$L$_\\odot$) star forming galaxy at the peak epoch of cosmic star formation. We find that MIPS J1428, contrary to average local ULIRGs, does not show a deficit in [O\\,I] relative to FIR. The combination of far-UV flux $G_0$ and gas density $n$ (derived from the PDR models), as well as the star formation efficiency (derived from CO and FIR) is similar to normal or starburst galaxies, despite the high infrared luminosity of this system. In contrast, F10214 has stringent upper limits on [O\\,IV] and [S\\,III], and an [O\\,III]/FIR ratio at least an order of magnitude lower than local starbursts or AGN, similar to local ULIRGs. } ",
        "introduction": "{\\it ISO} and {\\it Spitzer} observations have shown that the fraction of cosmic star formation in infrared luminous and ultra-luminous galaxies ( LIRGs/ULIRGs: L$_{IR}$=10$^{11}$L$_\\odot$ to a few times 10$^{12}$L$_\\odot$) increases from a few percent at z$\\sim$0 to more than 50\\% at z$>$1 (Genzel \\& Cesarsky 2000, Soifer et al. 2008). {\\it Herschel} offers the first opportunity to study the population of high-redshift infrared bright galaxies at wavelengths where they emit most strongly, both through photometry and high-resolution spectroscopy. In an attempt to probe for evolution of the ISM in the infrared galaxy population and to explore the limits of the PACS spectrometer we are investing several tens of hours for extremely deep high-resolution FIR spectroscopy of a (small) sample of galaxies with L$_{IR}>$10$^{12}$L$_\\odot$ at redshifts $\\ge$ 1. These data will be complemented by and compared to the results from our ongoing {\\it Herschel} spectroscopy of nearby ULIRGs, metal-rich and metal-poor galaxies, and starburst and AGN template objects. We will thus explore a wide range of parameter space and redshift. In this letter we present the first results on two high redshift objects, demonstrating the capabilities of the PACS spectrometer in long integrations and high-z spectroscopic observations. ",
        "conclusions": "We have detected [O\\,I]63$\\mu$m and [O\\,III]52$\\mu$m in MIPS J1428, and tentatively [O\\,III]52$\\mu$m in F10214. The spectra are shown in Figure \\ref{Fig1} (after re-binning to approximately the spectral resolution element, which is appropriate for resolved lines). These are the first detections of these lines in galaxies at such redshifts, i.e. at the epoch of the peak of cosmic star formation. Flux values, upper limits and continuum flux densities are given in Table \\ref{table:1}. Flux uncertainties from the calibration and from uncertainties in the continuum definition and line shapes are of the same order, and we estimate a total flux uncertainty of 40-50\\%.   {\\bf F10214:} we did not detect the targeted lines of [S\\,III]33$\\mu$m and [O\\,IV]26$\\mu$m, but [O III]52$\\mu$m (which is observed in the 1st order PACS band in parallel with the 2nd order [O IV] data) is serendipitously detected at a $\\sim$3$\\sigma$-level. We resolve the tentative line and measure a FWHM of 300$\\pm$100km/s (corrected for the instrumental profile, which has a FWHM$\\sim$220km/s at 170$\\mu$m), similar to the average CO line width (246$\\pm$10 km/s, Ao et al. 2008). Figure \\ref{Fig2} shows the ratio of the limit on [O\\,IV] to FIR\\footnote{FIR(42-122$\\mu$m)=1.26$\\times$10$^{-14}$(2.58$\\times f60 + f100$) [Wm$^{-2}$].} compared to a collection of local template objects from the literature (Graci\\'a-Carpio et al., in prep.). The FIR luminosity (L$_{FIR}$[40-500$\\mu$m]) has been calculated integrating SED data from the literature (Ao et al. 2008 and references therein) and our PACS continuum measurements, with a lens magnification correction factor of 12. The upper limit on [O\\,IV] is surprisingly low given that the source is classified as an AGN in the optical and mid-IR. Teplitz et al. (2006) report a detection of the [Ne\\,VI]7.6$\\mu$m line in their {\\it Spitzer}-IRS spectrum. Assuming an average [Ne\\,VI]/[O\\,IV] ratio of 0.4-0.6 in Seyfert galaxies (Sturm et al. 2002) and AGN-ULIRGs (Veilleux et al. 2009) we would have expected a two to three times higher line flux in [O\\,IV] than our upper limit. Only a few objects have a [Ne\\,VI]/[O\\,IV] ratio larger than 1. Given the uncertainties in both the {\\it Spitzer} and {\\it Herschel} measurements our non-detection might still be consistent with a [Ne\\,VI] detection. Figure \\ref{Fig2}, in any case, shows that the upper limit on the [O\\,IV]/FIR ratio in F10214 is lower than in Seyfert galaxies, and more in the regime of ULIRGs and starbursts. If the differential magnification model (see section \\ref{sect:targets}) is applied, then the upper limit on [O\\,IV] is yet lower by a factor of $\\sim$3. The low [O\\,IV]/FIR ratio in F10214 is consistent with L$_{FIR}$  arising predominantly from the starburst rather than the AGN, an assumption that is supported by model fits to the FIR SED, as F10214 has strong emission from the FIR to the (sub-)mm.  The [O\\,III]/FIR ratio for F10214 is shown in Figure \\ref{Fig3}b. It is at least an order of magnitude lower than local starbursts or AGN. There are not many ULIRG data points available for comparison. We have extracted two upper limits on \\object{Arp\\,220} and \\object{Mrk\\,231}, which are shown in the figure, from ISO-LWS data. Initial results from SHINING (Fischer et al. 2010) indicate a general deficiency of FIR fine-structure lines (including [O\\,III]) in local ULIRGs, similar to the [O\\,III] deficiency seen in F10214. With future data from SHINING we will be able to populate this diagram with many more data points for local templates covering a large parameter space, in order to put the findings reported here into context.   \\begin{figure} \\centering \\includegraphics[width=7cm]{14560fg2.eps} \\caption{The [O\\,IV]/FIR limit in F10214 compared to template objects compiled from the literature (Graci\\'a-Carpio et al.,  in prep.). Black symbols: AGN and ULIRGs known to harbour an AGN; blue: HII galaxies; green: LINERs. Open symbols and arrows are upper line flux limits. } \\label{Fig2}% \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=7cm]{14560fg3.eps} \\caption{[O I]63$\\mu$m and [O III]52$\\mu$m line strength in F10214 and MIPS J14284 as function of star formation efficiency. Comparison data are from the literature (Graci\\'a-Carpio et al., in prep.), except for [O\\,I] in Mrk\\,231 (Fischer et al. 2010). Symbols as in Figure \\ref{Fig2}.} \\label{Fig3}% \\end{figure}    {\\bf MIPS J1428}: the [O\\,I] and [O\\,III] lines (Figure \\ref{Fig1}) are detected with $\\sim$5 $\\sigma$ significance. They are redshifted by 220$\\pm$100 km/s and 300$\\pm$135 km/s, respectively, with respect to z=1.325 (derived from H$\\alpha$ and CO). Both lines are resolved with a FWHM of 750$\\pm150$km/s each (corrected for the instrumental profile). This is considerably broader than CO (380$\\pm$100 km/s, Iono et al. 2006), but comparable to H$\\alpha$ (530$\\pm$160km/s, Borys et al. 2006) within the errors. The measured continuum flux densities at 120 and 150 $\\mu$m (Tab. \\ref{table:1}) are low compared to the MIPS 160 $\\mu$m flux (430$\\pm$90mJy, Borys et al. 2006). However, preliminary processing of PACS and SPIRE photometry of the source (S. Oliver and H. Aussel, private communication) is consistent with our spectroscopic values. Using the PACS spectrum we derive L$_{FIR}$(40-500$\\mu$m)$\\sim$1.3$\\times$10$^{13}$L$_\\odot$/$\\mu$, where $\\mu$ is the magnification factor (i.e. L$_{FIR}\\ge$1.6$\\times$10$^{12}$L$_\\odot$ for $\\mu\\le$8). Contrary to F10214 the [O\\,III]/FIR ratio in MIPS J1428 (Figure \\ref{Fig3}) is of the same order as in local star forming and AGN galaxies. The [O\\,III]/[O\\,I] ratio is very similar to the ratio in typical starburst galaxies (like M\\,82).   For the [O\\,I] spectrum we cannot rule out an underlying absorption of ammonia (NH$_3$). In Figure \\ref{Fig1}, top panel, the dotted line is a NH$_3$ model spectrum (arbitrarily scaled). The line at the bottom of panel (a) is the residual after subtracting the line (+continuum) fit and the NH$_3$ model from the data. Strong FIR NH$_3$ features with high column densities have been detected before in infrared bright galaxies, in the infrared (e.g. in Arp\\,220, Gonz\\'alez-Alfonso et al. 2004) and at mm wavelengths (e.g. Henkel et al. 2008). On the other hand, the corresponding NH$_3$ column density in MIPS J1428 would be quite high, and the object is not known to be heavily obscured at other wavelengths.    [O\\,I] and [C\\,II] are usually the brightest cooling lines of the cool ISM in galaxies. However, ISO data have shown the [C\\,II] line in local ULIRGs to be about an order of magnitude lower relative to the FIR continuum than in normal and starburst galaxies (e.g. Malhotra et al. 2001, Luhman et al. 2003). The relative weakness of [C\\,II] has strong implications on its potential use as a star formation tracer for high redshift (z$\\ge$4) studies (e.g. Maiolino et al. 2005 and 2009). Hailey-Dunsheath et al. (2010) have observed [C II] with the ZEUS spectrometer at the CSO. They found MIPS J1428 to be a counter example, with a normal [C\\,II]/FIR ratio. Our [O\\,I] observation (Fig. \\ref{Fig3}) supports this, i.e. this source, a high redshift source with the high SFR and L$_{IR}$ of a local ULIRG, does not show a deficit in the major PDR cooling lines relative to the FIR.  We plot in Figure \\ref{Fig5} the position of MIPS J1428 in a [C\\,II]/[O\\,I] versus ([C\\,II]+[O\\,I])/FIR diagram. For reference, we overplot the PDR model curves from Kaufman et al. (1999), and show the location of various samples of template galaxies from the literature. This is the first time that a combined [C\\,II] and [O\\,I] PDR diagnostic is possible for a source at redshift $\\sim$1. The PDR model curves are appropriate for clouds in the active regions of galaxies that are illuminated on all sides (see Contursi et al. 2002). Combined HII region and PDR modeling of normal and starbursting galaxies have demonstrated that PDRs account for more than half of the [C\\,II] emission in these sources (e.g. Colbert et al. 1999). Strong shocks could boost the [O\\,I] emission relative to [C\\,II]. However, the measured [C\\,II]/[O\\,I] ratio of 2.5$\\pm$1.2 is well above 1, suggesting we can ignore the role of shocks in this case. We conclude that local ULIRGs and AGNs are not good analogs of MIPS J1428, and that this source appears to be more like a normal/starburst galaxy.  Typical massive galaxies in the distant Universe formed stars an order of magnitude more rapidly than in the local Universe. Either star formation was significantly more efficient or these young galaxies were much more gas rich (Tacconi et al. 2010). Figure \\ref{Fig3} shows the [O\\,III]/FIR ratio as a function of star formation efficiency (SFE, expressed in L$_{FIR}$/M$_{gas}$). M$_{gas}$ was calculated from CO measurements (Iono et al. 2006, Ao et al. 2008). For F10214 this assumes a CO conversion factor like in local ULIRGs, and  that the FIR arises from star formation rather than an AGN, an assumption that is supported by the low [O\\,IV]/FIR ratio and model fits to the FIR SED (see above). For MIPS J1428 we used a CO conversion factor for normal star forming galaxies, motivated by our findings above, which yields M$_{gas}$ = 5$\\times$10$^{11}$M$_\\odot$/$\\mu$. The star formation rate in MIPS J1428 of 2600M$_\\odot$/yr/$\\mu$ (i.e. SFR $\\gtrsim$ 300M$_\\odot$/yr for $\\mu\\le$ 8), is quite high, comparable to local ULIRGs or to SMGs. However, given the larger gas reservoir, the star formation {\\it efficiency} is well in the range of normal star forming galaxies. Even with a ULIRG-like conversion factor, L$_{FIR}$/M$_{gas}$ would still be lower than in most local ULIRGs. Irrespective of the conversion factor we find L$_{FIR}$/L$_{CO}$ = 115L$_\\odot$/(K km/s pc$^2$) for MIPS J1428. In a recent analysis of CO data from z$\\sim$1-3 normal star forming galaxies Genzel et al. (2010) find values of 5 to 100 for this ratio, while the average ratio for major mergers is $\\sim$160, reaching values up to 600. This is consistent with our interpretation of MIPS J1428 as a normal star forming galaxy. Our result thus adds further support to the recent results from the SINS survey (Tacconi et al. 2010) which indicate that with an increased gas reservoir star forming galaxies at high redshifts can achieve ULIRG luminosities without being major mergers.   Upcoming SHINING and other {\\it Herschel} observations will enable us to compare with more template objects covering a larger parameter space and to search for galaxy evolution effects in the ISM of infrared bright galaxies.      \\begin{figure} \\centering \\includegraphics[width=9.3cm]{14560fg4.ps} \\caption{ [C\\,II]/[O\\,I] vs. ([C\\,II]+[O\\,I])/FIR. PDR models of Kaufman et al. (1999) are used. The position of MIPS J1428 is indicated, using [C II] from Hailey-Dunsheath et al. (2010). Symbols as in Figure \\ref{Fig2}, sources with L$_{IR}\\ge$10$^{12}$L$_\\odot$ are marked with a cross. Open symbols correspond to [O\\,I] non-detections. The Mrk 231 data are from Fischer et al. (2010).} \\label{Fig5}% \\end{figure}"
    },
    "1005/1005.4933_arXiv.txt": {
        "abstract": "PKS 0558\u00ad-504 is the brightest radio-loud Narrow-Line Seyfert 1 galaxy at X-ray energies. Here we present results from the radio, optical, UV, and X-ray bands obtained with \\swift, \\xmm, and \\atca\\ during a 10-day monitoring campaign in September 2008. The simultaneous coverage at several wavelengths makes it possible to investigate in detail the broadband spectral energy distribution (SED) and the energetic of this source. The main results can be summarized as follows. The \\atca\\ reveals the presence of an extended radio emission in PKS~0558--504 with two lobe-like structures at $\\sim$7\\arcsec\\ from the bright central source. The extended radio structure and the low value of the radio-loudness similar to radio-quiet Seyfert galaxies coupled with constraints from higher energy bands argue against a jet-dominated emission in \\pks. The study of the SED, which is dominated by a nearly constant optical-UV emission, supports the conclusion that PKS~0558--504 is accreting at super-Eddington rate. This conclusion was reached assuming $M_{\\rm BH}=2.5\\times10^8~{M_\\odot}$, which was obtained with a new scaling method based on X-ray spectral variability results. A comparison between the accretion luminosity and the kinetic power associated with the jet suggests that in this source the accretion power dominates in agreement with the results obtained from Radiation-MHD simulations of Galactic black holes (GBHs) accreting at the Eddington rate. The combined findings from this panchromatic investigation strongly suggest that \\pks\\ is a large scale analog of GBHs in their highly accreting intermediate state. Importantly, \\pks\\ may also be the prototype of the parent population of the very radio-loud NLS1s recently detected at $\\gamma$-ray energies. ",
        "introduction": "Bipolar relativistic jets are common features in a variety of astrophysical objects, most notably in Galactic Black Holes (GBHs) and Active Galactic Nuclei (AGNs). Accretion of gas onto black holes is thought to power these collimated outflows. However, the details of the jet formation as well as the nature of the coupling between accreting matter and outflows are still among the outstanding open questions in high energy astrophysics.  Because of their vicinity and hence their high brightness, the temporal and spectral properties of GBHs are much better known and can be used to infer information on their more powerful, extragalactic analogs. Indeed considerable progress in this field has been made by multi-wavelength correlated studies of GBHs in different spectral states, which allow one to study the link between accretion (generally probed by the X-ray emission) and jet properties (in the radio regime) on ``human'' timescales. It is now well established that GBHs undergo state transitions, switching between two main states: the low/hard (LS) and the high/soft state (HS) passing through  soft and hard intermediate states (IS), which are also named very high states (VHS) when they occur at high values of accretion rate (see \\citealt{mcclin06,done07} for recent comprehensive reviews on GBHs). Each spectral state is unambiguously characterized by a specific combination of X-ray temporal and spectral properties and by well defined radio features \\citep[see][]{fend04}.  In the study of the disk-jet link one of the most interesting spectral states is the IS/VHS, which is generally characterized by powerful transient relativistic ejections in the radio coupled with highly variable X-ray emission that is unambiguously associated with the accretion flow, unlike the LS where the origin of the X-rays is still matter of debate \\citep[e.g.][]{mark03,zdzi04}.  Unfortunately, the physical conditions that lead to the transient relativistic ejections during the IS/VHS are still poorly understood mostly because of their short duration in GBHs. Since the dynamical time scales are proportional to the black hole mass, in individual AGNs it is not possible to observe long-term phenomena occurring in GBHs such as the canonical spectral transitions. On the other hand, AGNs may provide better constraints on short-lived GBH phenomena and hence shed light on the jet formation and the interplay between accretion and ejection processes. \\begin{table*}[bht] \\scriptsize \\caption{\\swift\\ Observation log of PKS 0558-504} \\begin{center} \\begin{tabular}{lcccrrrrrrr} \\hline \\hline \\noalign{\\smallskip} Segment & Start time  & End Time & MJD & \\multicolumn{7}{c}{Observing time given in s}\\\\ &(UT) &(UT) & & $\\rm T_{\\rm XRT}$& $\\rm T_{\\rm V}$ & $\\rm T_{\\rm B}$ &$\\rm T_{\\rm U}$ & $\\rm T_{\\rm UVW1}$ & $\\rm T_{\\rm UVM2}$ & $\\rm T_{\\rm UVW2}$\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 001 & 2008-09-07 08:10 & 2008-09-07 11:36 & 54716.41 & 2215 & 122 & 122 & 122 & 244 & 210 & 488 \\\\ 002 & 2008-09-08 01:50 & 2008-09-08 06:50 & 54717.20 & 2429 & 199 & 199 & 199 & 399 & 358 & 798 \\\\ 003 & 2008-09-09 09:52 & 2008-09-09 11:48 & 54718.45 & 2327 & 191 & 191 & 191 & 381 & 532 & 763 \\\\ 004 & 2008-09-10 00:15 & 2008-09-10 18:20 & 54719.39 & 2203 & 181 & 181 & 181 & 361 & 509 & 724 \\\\ 005 & 2008-09-11 13:21 & 2008-09-11 19:59 & 54720.69 & 1897 & 163 & 163 & 163 & 326 & 344 & 653 \\\\ 006 & 2008-09-12 00:36 & 2008-09-12 05:37 & 54721.15 & 2008 & 155 & 155 & 155 & 310 & 483 & 621 \\\\ 007 & 2008-09-13 15:09 & 2008-09-13 18:40 & 54722.68 & 1869 & 148 & 148 & 148 & 296 & 447 & 594 \\\\ 008 & 2008-09-14 00:42 & 2008-09-14 05:40 & 54723.13 & 2363 & 240 & 240 & 240 & 483 & 682 & 966 \\\\ 009 & 2008-09-15 16:56 & 2008-09-15 21:54 & 54724.80 & 2353 & 184 & 184 & 184 & 369 & 557 & 739 \\\\ 010 & 2008-09-16 07:20 & 2008-09-16 12:11 & 54725.90 & 2072 & 165 & 165 & 165 & 330 & 461 & 660 \\\\ \\noalign{\\smallskip} \\hline \\hline \\end{tabular} \\end{center} \\label{tab1} \\end{table*}  In the framework of the AGN-GBH unification, the Narrow-Line Seyfert 1 galaxies (NLS1s) are the best candidates for large-scale analogs of GBHs in the IS/VHS. NLS1s are historically identified by their optical emission line properties: the ratio [O III]/H$\\beta$ is less than 3 and FWHM H$\\beta$ is less than 2000${~\\rm km~s^{-1}}$ \\citep{oster85,good89}. They are seldom radio loud \\citep{komo06}, although recent studies reveal the existence of several NLS1s characterized by very high radio-loudness \\citep[e.g.][]{yuan08}. Recently, a few of these very radio-loud NLS1s have been detected at $\\gamma$-ray energies by the {\\it Fermi/LAT} collaboration confirming that these sources posses relativistic jets observed at small viewing angles \\citep{abdo09a,abdo09b,abdo09c, fosc09}. In the X-rays, NLS1 are generally characterized by steep spectra and strong variability \\citep[e.g.][]{brand99,lei99a,lei99b,gru01}. Based on these properties, it has been suggested that NLS1 are AGN in their early phase \\citep{gru99,math00}, characterized by relatively small black hole masses \\citep[e.g.][]{gru04}, and very high accretion rates in terms of Eddington units (e.g. \\citealt{boro92,sule00,gru10}; see also \\citealt{marco08,deca08} for a discording view).  PKS~0558--504 ($z=0.137$) is the brightest radio-loud ($R=L_\\nu(6~{\\rm cm})/L_\\nu({\\rm B})\\simeq 27$, \\citealt{sieb99}) NLS1 in the X-ray band and therefore one of the best studied at these energies. Previous X-ray observations with different satellites have confirmed that \\pks\\ shows the characteristic NLS1 properties: strong variability, steep X-ray spectrum, substantial soft excess, and relatively high luminosity. By comparing the X-ray observations from several satellites over a decade, it is evident that on long timescales the strong X-ray variability of PKS~0558--504 occurs persistently but without being accompanied by significant spectral variability: the spectrum above 2 keV is consistently described by a power-law model with photon index  $\\Gamma\\sim2.2$ \\citep{glioz00}. This conclusion  has been confirmed by a recent monitoring campaign with \\rxte: the long-term achromatic variability appears to be consistent with the behavior of GBHs during the transition LS-to-HS and inconsistent with the spectral variability typical of jet-dominated AGNs \\citep{glioz07}.  In the last ten years, PKS~0558--504 has been observed repeatedly by \\xmm\\ \\citep{obri01,glioz01,brink04}. The most recent and highest quality data, obtained with the \\xmm\\ EPIC cameras from five consecutive orbits in September 2008, confirm that \\pks\\ is highly variable on all sampled timescales, and a spectral analysis indicates that the 2--10 keV energy range is well fitted with a power law whose slope steepens as the source brightens and the soft excess is well described by a low-temperature Comptonization model \\citep{papa10a}. These findings suggest that despite its radio loudness \\pks\\ appears to behave like a ``normal'' Seyfert galaxy from the X-ray point of view.   The main goal of this work is to shed light on the energetic of this powerful source using multi-wavelength data. For this purpose, after describing the observations and data reduction in $\\S~2$, we first try to assess the role played by the jet emission beyond the radio band (in $\\S~3$) and then to constrain the black hole mass in \\pks\\ without the use of optical measurements (in $\\S~4$). In Section 5 we describe the multi-wavelength behavior of \\pks\\ based on a 10-day multi-wavelength campaign carried out with the \\swift\\ XRT and UVOT simultaneously with the deep \\xmm\\ observations in September 2008, and complemented with one {\\it radio} observation with the Australia Telescope Compact Array (ATCA). Finally, in $\\S~6$ we discuss the implications of the main results and draw our conclusions. Hereafter, we adopt a cosmology with $H_0=71{\\rm~km~s^{-1}~Mpc^{-1}}$, $\\Omega_\\Lambda=0.73$ and $\\Omega_{\\rm M}=0.27$ \\citep{ben03}; with the assumed cosmological parameters, the luminosity distance of \\pks\\  is 642 Mpc. ",
        "conclusions": "\\subsection{Jet contribution} The main goal of this work is to investigate the energetic of the radio-loud NLS1 \\pks\\ and shed some light on the nature of its central engine. To this end, it is first necessary to assess the role played by the jet and more specifically whether the jet may dominate the high energy emission of \\pks. The findings based on the \\atca\\ observations reported in $\\S3$ -- the radio-loudness measurements consistent with radio-quiet Seyfert galaxies and the extended and symmetric FRI-like radio structure which is at odds with the compact radio emission typical of very radio-loud NLS1s-- disfavor a jet-dominated scenario, which instead appears to be the most likely scenario for the very radio-loud NLS1s recently detected with the {\\it Fermi/LAT}.  This preliminary conclusion is supported by several independent pieces of evidence from studies at higher energies, which can be summarized as follows: (1) the X-ray spectral and temporal variability appears to be fully consistent with the typical behavior of radio-quiet Seyfert galaxies \\citep{papa10a,papa10b}; (2) the long-term spectral variability in the 2--15 keV energy band is inconsistent with the characteristic trend observed in blazars with \\rxte\\ \\citep{glioz07}; (3) unlike highly beamed AGNs, \\pks\\ is not detected at very high energy neither in the GeV range (the first year Fermi LAT catalog) nor at TeV energies (HESS; Giebels private communication). Combining all the information available for \\pks\\ it is thus reasonable to conclude that the jet emission does not play a relevant role beyond the radio band.  Therefore, for the rest of the paper we will assume that the bulk of the optical-UV-X-ray emission is associate with the accretion flow only.    \\subsection{Energetic of \\pks} The main result from the analysis of the broadband SED  (described in $\\S5$) is that the emission of \\pks\\ is dominated by a weakly variable UV bump ($F_{\\rm var,UV}\\simeq 4\\%$), which is commonly associated with the emission from the accretion disk. Conversely, the X-ray radiation, which we have interpreted as emission from two Comptonization components, appears to be highly variable ($F_{\\rm var,X}\\simeq 30\\%$) but encompasses only a small fraction of the total emission. It is instructive to assess in a more quantitative way which of the variable components (UV or X-rays) is energetically more important. Since $F_{\\rm var}$ provides a measurement of the variability normalized over the average flux, a measurement of the ``true\" variability (hereafter $var_{\\rm UV}$ and $var_{\\rm X}$, respectively) can be obtained by multiplying $F_{\\rm var}$ by the respective average flux. This test yields $var_{\\rm X}/var_{\\rm UV}=1.3$, indicating that the X-ray component is the dominant variable component.  In order to get helpful insights into the nature of the central engine in \\pks\\ it is crucial to constrain its accretion rate in term of Eddington units. This information can be readily obtained by knowing the mass of the central black hole and measuring the bolometric luminosity.  Depending on the method used (see $\\S4$) the black hole mass of \\pks\\  may range between $6\\times10^7{\\rm ~M_\\odot}$ and  $3\\times10^8{\\rm ~M_\\odot}$. To be conservative, in the rest of the paper we will assume $M_{\\rm BH}=2.5\\times 10^8{\\rm ~M_\\odot}$ which corresponds to an Eddington luminosity of $L_{\\rm Edd}=3.25\\times 10^{46}$ \\lum\\ (note that this choice of $M_{\\rm BH}$ does not affect the main findings nor our conclusions). Combining this value with the average bolometric luminosity derived from the direct integration of the broadband SED (see Table 3) leads to an accretion rate of $\\lambda_{\\rm Edd}\\equiv L_{\\rm bol}/L_{\\rm Edd}\\simeq 1.7$ that strongly suggests that \\pks\\ accretes at super-Eddington rate. Note that with the less conservative assumption of $M_{\\rm BH}=6\\times 10^7{\\rm ~M_\\odot}$ (which corresponds to a lower Eddington luminosity and a higher bolometric luminosity), the accretion rate in Eddington units would be of the order of $\\lambda_{\\rm Edd}\\simeq 25$.  Having measured the X-ray and bolometric luminosity over a period of ten days, it is possible to compute not only the average bolometric correction $\\kappa_{\\rm 2-10~keV} \\equiv L_{\\rm 2-10~keV}/L_{\\rm bol}$ but also to investigate if it varies over time. From Table 3 we infer that $\\kappa_{\\rm 2-10~keV}$ indeed varies by a factor of $\\sim$2.5 ranging between 108 and 249, with an average value of $162\\pm39$ (where the error quoted is 1$\\sigma$). The large value of $\\kappa_{\\rm 2-10~keV}$ appears to be in qualitative agreement with recent studies that suggest that the X-ray bolometric correction is directly proportional to the Eddington ratio \\citep[e.g.][]{vasu07,vasu09}. A more quantitative comparison can be performed using the recent results presented by \\citet{lusso10}, which are based on the optical and X-ray study of a sample of 545 X-ray selected type 1 AGNs from the XMM-COSMOS survey. More specifically, plugging $\\lambda_{\\rm Edd}=1.7$  (the value of the Eddington ratio derived for \\pks) into equation (14) from \\citet{lusso10} we derive $\\kappa_{\\rm 2-10~keV}=151_{-20}^{+24}$, which is fully consistent with the average value empirically derived from the SED.  Since \\pks\\ possesses a radio jet, it is instructive to estimate the the inflow accretion rate and try to compare it to the the outflow mass rate associated with the jet, and similarly compare the accretion luminosity to the kinetic power. With the reasonable assumption that $L_{\\rm bol}=L_{\\rm accr}=\\eta\\dot M_{\\rm in} c^2$, we derive an accretion inflow rate of $\\dot M_{\\rm in} \\simeq 1 {~\\rm M_\\odot/yr}$, assuming a typical radiative efficiency of $\\eta=0.1$ (larger values of $\\dot M$ are obtained if photon trapping effects play a significant role causing $\\eta$ to decrease considerably). Following \\citet{merlo07} it is possible to derive an estimate of the kinetic power of the jet for radio galaxies by making use of the formula $\\log L_{\\rm kin}=0.81\\log L_{\\rm R} +11.9$, where $L_{\\rm R}$ is the nuclear radio luminosity at 5 GHz. Using the  value measured with the \\atca, $L_{\\rm 5 GHz}=1.7\\times10^{41}$ \\lum, we  derive a kinetic power of the order of $L_{\\rm kin}\\simeq 2\\times10^{45}$ \\lum. The latter is considerably lower than the luminosity associated with the accretion flow, which is of the order of $5\\times10^{46}$ \\lum. We can estimate the outflow mass rate from $L_{\\rm kin}=(\\Gamma-1)\\dot M_{\\rm out}c^2$ by assuming an appropriate value for the bulk Lorentz factor. With $\\Gamma=3$, which is a typical value for radio galaxies and consistent with the value inferred from the brightness temperature derived from the radio variability (Kedziora-Chudczer et al. 2010, in preparation), we derive $\\dot M_{\\rm out} \\simeq 2\\times10^{-2} {~\\rm M_\\odot/yr}$, which is about two orders of magnitude lower than the value inferred for the accretion rate $\\dot M_{\\rm in}$.    \\subsection{Analogy with GBHs} Our recent study based on a 1-year \\rxte\\ monitoring campaign suggests that the spectral variability behavior of \\pks\\ in the 2--15 keV energy band mimics GBHs in their highly accreting IS \\citep{glioz07}. We can now test this analogy in a more quantitative way. First, we can compare the results from the energetic section relative to the disk-jet interaction with the findings of \\citet{osu09} that are derived from radiation-magneto-hydrodynamic simulations of a stellar BH system accreting at the Eddington rate. The general agreement with the simulations (specifically, $\\dot M_{\\rm out}\\ll \\dot M_{\\rm in}$ and $L_{\\rm kin} < L_{\\rm accr}$) lends further support to the hypothesis that \\pks\\ is a large scale analog of a GBH in a highly accreting state.  The analogy with the IS/VHS is supported by the direct comparison of \\pks\\ with GRO J1655-40 in the $\\Gamma - N_{\\rm BMC}$ diagram that we used to estimate the black hole mass (see Fig.~\\ref{figure:fig3}). Specifically, the common spectral variability trend occurs when GRO J1655-40 is switching from the LS to the HS (i.e., the source is indeed in the IS) during the rise of the outburst in March 2005 (see details in \\citealt{shap07}). Interestingly, the X-ray outburst of GRO J1655-40 was accompanied by enhanced emission in the radio band that faded before the peak of the X-ray outburst. Based on the values reported by \\citet{shap07}, the ratio between radio/X-ray and radio/optical fluxes places GRO J1655-40 in the same lower-left quadrant of the radio-loudness plot shown in Fig.~\\ref{figure:fig2} for \\pks.  Further support to this conclusion comes from the comparison of the broadband spectral index $\\alpha_{\\rm OX}$ of \\pks\\ with the corresponding quantity $\\alpha_{\\rm GBH}$, which was recently introduced by \\citet{sobo09} to characterize the broadband spectrum of GBHs and allows a direct comparison with AGNs. During our short multi-wavelength campaign in October 2008, $\\alpha_{\\rm OX}$ ranged between 1.25--1.4 while the photon index varied from 2.1--2.25. Importantly, this is the very range corresponding to the intermediate state in GBHs (see Figure 6 in \\citealt{sobo09}).  Finally, it is instructive to compare \\pks\\ to GRS 1915+105, which is thought to be the only microquasar that regularly accretes at or above the Eddington limit. Interestingly, during its spectral transition hard-to-soft the X-ray spectra of GRS 1915+105 are well fitted by two BMC components (a hard component with $\\Gamma_1=$1.7--3 and a soft component with $\\Gamma_1=$2.7--4.2) which evolve following the trend shown in Fig.~\\ref{figure:fig3} \\citep{tita09}. Also for GRS 1915+105 the location of the IS/VHS in the $\\Gamma - N_{\\rm BMC}$ diagram is consistent with the position of \\pks\\ (see Figure 13 of \\citealt{tita09}).  \\subsection{Conclusion} In conclusion, our multiwavelength analysis of the emission properties of \\pks\\ ranging between the radio and the hard X-rays reveals that this source has a radio jet, yet it  is accreting at the Eddington level or above. All our findings appear to independently lend support to the hypothesis that \\pks\\ may be a large scale analog of the IS observed in GBHs. But can we extend this analogy to the whole class of NLS1s? In other words, can \\pks\\ be considered a prototype of NLS1s? At first sight, the ``radio-loudness\" of \\pks\\ seems to set this source in a small sub-class of radio-loud NLS1. However, $R=10$ does not represent a strict boundary for the radio-loudness; indeed the radio-loudness of \\pks\\ appears to be fully consistent with that of radio-quiet Seyfert galaxies according to both $R$ and $R_x$ (see Fig.~\\ref{figure:fig2}). Bearing in mind that \\pks\\ is by far the brightest NLS1 in the X-ray band and that its X-ray luminosity is nearly 100 times the typical value of ``normal\" radio-quiet NLS1s, similar values of the order of $R_x \\la 10^{-3}$ for ``normal\" NLS1s would correspond to radio fluxes of the order of 1 mJy or less. This would explain why only a small minority of NLS1s is detected in the radio band. Additionally, it must be pointed out that in the IS the radio ejections are transients, therefore, genuine radio-quiet NLS1 might still represent large scale analogs of the IS.  It is also interesting to compare \\pks\\ with the radio-loud NLS1s recently detected at $\\gamma$-rays. The latter are characterized by large values of the radio-loudness ($R > 50$), by compact radio emission, and by SEDs that are dominated by two broad bumps typical of jet-dominated sources. Indeed, it has been argued that these very radio-loud NLS1s galaxies represent a the third class of $\\gamma$-ray emitting AGNs along with blazars and radio galaxies. In this perspective, \\pks\\ can play a crucial role since it may represent the prototype of NLS1s with a jet seen at large viewing angles. In other words, \\pks\\ may be the the first representative of non-jet-dominated radio-loud NLS1s which are the parent population of the very radio-loud NLS1s detected by the {\\it Fermi/LAT}, similarly to radio galaxies that are considered the non-beamed parent population of blazars.   In summary, we can hypothesize that \\pks\\ is indeed a prototypical NLS1. Since the jet appears to dominate  only in the radio band, the emission associated with the accretion flow can be cleanly separated by that related to the jet. This makes \\pks\\ an ideal system to study the disk-jet interaction. In this framework, important insights into the physical conditions leading to the formation of jets in highly accreting BH systems can be obtained by a systematic comparison of the spectral and temporal properties of \\pks\\ with the corresponding properties of other highly accreting NLS1 that do not show any evidence of radio jets as well as with NLS1s that are jet-dominated. We plan to pursue this approach in our future work."
    },
    "1005/1005.3208_arXiv.txt": {
        "abstract": "{Accurate parameters of the host stars of exoplanets are important for the interpretation of the new planet systems that continue to emerge. The \\corot\\ satellite recently discovered a transiting rocky planet with a density similar to the inner planets in our solar system, a so-called \\SE. The mass was determined using ground-based follow-up spectroscopy, which also revealed a second, non-transiting \\SE.} {These planets are orbiting a relatively faint ($m_V=11.7$) G9V star called CoRoT-7. We wish to refine the determination of the physical properties of the host star, which are important for the interpretation of the properties of the planet system. } {We have used high-quality spectra from HARPS@ESO\\,3.6m and UVES@VLT\\,8.2m. We use various methods to analyse the spectra using 1D~LTE atmospheric models. From the analysis of \\feone\\ and \\fetwo\\ lines we determine the effective temperature, surface gravity and microturbulence. We use the Balmer lines to constrain the effective temperature and pressure sensitive Mg\\,1b and Ca lines to constrain the surface gravity. We analyse both single spectra and co-added spectra to identify systematic errors. We determine the projected rotational velocity and macroturbulence by fitting the line shapes of isolated lines. We finally employ the Wilson-Bappu effect to determine the approximate absolute magnitude.} {From the analysis of the three best spectra using several methods we find $T_{\\rm eff}=5250\\pm60$\\,K, $\\log g = 4.47\\pm0.05$, $[{\\rm M/H}]=+0.12\\pm0.06$, and $v \\sin i = 1.1^{+1.0}_{-0.5}$\\,\\kms. The chemical composition of $20$ analysed elements are consistent with a uniform scaling by the metallicity $+0.12$~dex. We compared the $L/M$ ratio with isochrones to constrain the evolutionary status. Using the age estimate of 1.2--2.3~Gyr based on stellar activity, we determine the mass and radius $0.91\\pm0.03$\\,\\msun\\ and $0.82\\pm0.04$\\,\\rsun. With these updated constraints we fitted the \\corot\\ transit light curve for \\corotp. We revise the planet radius to be slightly smaller, $R = 1.58\\pm0.10$~R$_\\oplus$, and using the planet mass the density becomes slightly higher, $\\rho = 7.2\\pm1.8 \\, {\\rm g\\,cm}^{-3}$.  } {The host star \\corots\\ is a slowly rotating, metal rich, unevolved type G9V star. The star is relatively faint and its fundamental parameters can only be determined through indirect methods. Our methods rely on detailed spectral analyses that in turn depend on the adopted model atmospheres. From the analysis of spectra of stars with well-known parameters with similar parameters to \\corots\\ (the Sun and \\acenb) we demonstrate that our methods are robust within the claimed uncertainties. Therefore our methods can be reliably used in subsequent analyses of similar exoplanet host stars. } ",
        "introduction": "\\begin{figure*} \\centering \\includegraphics[width=17cm,angle=0]{bruntt_f01a.eps} % \\includegraphics[width=17cm,angle=0]{bruntt_f01b.eps} % \\caption{A section of the CoRoT\\,7/H1-7 spectrum illustrating how the spectrum is normalised with \\rainbow. The top panels show a wide range and the lower panels show a zoom near the edge of the same echelle order. The neighbouring order is shown with a short-dashed line. The thick long-dashed line is a spline fit to the continuum windows marked by circles. The normalised spectrum agrees reasonably well with the template synthetic spectrum (smooth green line). The agreement between the two overlapping orders is good and will finally be merged to improve the S/N. \\label{fig:rainbow}} \\end{figure*} The discovery of the first \\SE~planet outside of the Solar system with a measured {\\em absolute} mass and radius was recently announced, based on photometric data from \\corot\\ (Convection, Rotation and planetary Transits; \\citealt{baglin06}). This planet, \\corotp, has a radius of $1.68\\pm0.09$ \\rearth\\ \\citep{leger09}, mass $4.8\\pm0.8$ \\mearth\\ \\citep{queloz09}, and the orbital period is about 0.85 days \\citep{leger09}. The average density is $5.6\\pm1.3$\\,g\\,\\cmthree\\ which is similar to Mercury, Venus and the Earth \\citep{queloz09}. Furthermore, a second non-transiting \\SE\\ has been found from radial-velocity monitoring \\citep{queloz09}. These results have only been possible to achieve thanks to an extensive ground-based follow-up program of the relatively faint star \\corots\\ (TYCHO~ID 4799-1733-1; $m_V = 11.7$) over more than one year.  In the derivation of the planetary parameters, one of the most important factors is the correct identification of the host star's fundamental parameters and evolutionary stage. It is particularly important to estimate the stellar radius which is imperative for determining the absolute planetary radius. A first analysis by several of the \\corot\\ teams has been carried out in \\cite{leger09}, based on a spectrum from the ``Ultraviolet and Visual Echelle Spectrograph'' (UVES@VLT\\,8.2m) and a preliminary analysis of 53 co-added spectra from the ``High Accuracy Radial velocity Planet Searcher'' (HARPS@ESO\\,3.6m).   Since then, a total of 107 spectra from \\harps\\ have become available \\citep{queloz09}. These spectra have higher spectral resolution and better signal-to-noise (S/N) than the UVES spectrum analysed by \\cite{leger09}. We can therefore now refine the analysis of \\corots\\ and possibly impose stronger constraints on the properties of the system. The methods we employ have been developed during the analysis of other \\corot\\ targets \\citep{deleuil08, moutou08, rauer09, fridlund10, bruntt09}. In the current paper we have expanded these tools and we will describe our approach in greater detail than previously done. These tools will be the standard methods to be applied for the characterization of future \\corot\\ targets. ",
        "conclusions": "We have presented a detailed spectroscopic analysis of the planet-hosting star \\corots. The analysis is based on \\harps\\ spectra which have higher signal-to-noise and better resolution than be \\uves\\ spectrum used to get a preliminary result \\citep{leger09}. We analysed both individual spectra from different nights and co-added spectra and found excellent agreement. Only for one of the single \\harps\\ spectra did we find a systematic error in \\teff\\ and \\logg, which is explained by the low S/N.  We described in detail the \\vwa\\ tool which is used for determination of the atmospheric parameters \\teff\\ and \\logg\\ using \\feonetwo\\ lines and the pressure sensitive \\mgoneb\\ and Ca lines. We used the \\sme\\ tool to analyse in addition the Balmer and \\naone\\ lines. We find excellent agreement between the different methods.  To evaluate the evolutionary status (age) and fundamental stellar parameters (mass, radius) we compared the observed properties of \\corots\\ with theoretical isochrones. From the spectroscopic \\teff\\ and \\logg\\ we can estimate the $L/M$ ratio. We compared this with isochrones but find that the uncertainty is too large to constrain the evolutionary status. However, by imposing constraints on the stellar age (1.2--2.3~Gyr from \\citealt{leger09}) we can constrain the mass and radius to $0.91\\pm0.03$\\,\\msun\\ and $0.82\\pm0.04$\\,\\rsun. This is a only slight revision of the original value from \\cite{leger09} who used a lower metallicity. % The relatively large uncertainty of 7\\% on the stellar radius directly impacts the accuracy of the determine radius and density of the transiting planet, \\corotp.  We have used the new stellar parameters to fit the transit light curve reported by \\cite{leger09}. We used the formalism of \\cite{gimenez06} with fixed limb-darkening coefficients, and we explored the parameter space which is consistent with the stellar parameters and their associated uncertainty. The constraints to the fit include the orbital inclination ($81.45\\pm1.10^\\circ$), the phase of transit ingress $\\theta_1 (0.02785\\pm0.00005$), and the ratio of planet-to-star radius ($0.0176\\pm0.0003$). We refer to Sect.~9 in \\cite{leger09} for a description of the methodology of the fitting procedure. With the new stellar parameters we determine the radius of the planet to be slightly smaller with radius $1.58\\pm0.10$~R$_\\oplus$ (\\citealt{leger09} found $1.68\\pm0.09$~R$_\\oplus$). The slightly smaller radius is mainly due to our revision of the stellar radius.  The new stellar mass and the updated inclination were used, together with the published values of the ephemeris \\citep{leger09} and radial velocity semiamplitude \\citep{queloz09} to estimate the mass of the planet CoRoT-7b as $5.2 \\pm0.8 \\, {\\rm M}_\\oplus$. Combined with the radius of the planet this results in a density of $7.2 \\pm 1.8 {\\rm \\,g\\,cm}^{-3}$. which is consistent, but slightly more dense than the reported value of $5.6\\pm1.3$ g\\,cm$^{-3}$ in the previous work.     We also analysed spectra of the Sun and \\acenb, also observed with the \\harps\\ spectrograph. For these stars the fundamental parameters are known with very good accuracy and they can therefore be used to validate the methods we use for the much fainter star \\corots. We compared the spectroscopically determined \\teff\\ and \\logg\\ with the values from fundamental methods for \\acenb, \\ie\\ using the binary dynamical mass and the interferometric determination of the radius. There is excellent agreement within 1-$\\sigma$, indicating that the adopted uncertainties are realistic. This gives us some confidence that we can use theoretical evolution models to constrain the radii and masses of stars, but requires that limits can be put on the stellar age.  The exoplanet host star \\corots\\ is a slowly rotating, metal rich, type G9V star. The star is relatively faint and its fundamental parameters can only be determined through indirect methods. The expected future discoveries of similar planet systems with \\corot\\ and \\kepler\\ will also be limited by our ability to characterise the host stars. In the case of \\kepler\\ we have the additional advantage that for the brightest stars the solar-like pulsations can be used to constrain the stellar radius \\citep{jcd10}. This analysis also relies on evolution models but will be able to constrain the stellar radius to about 2\\% \\citep{stello09, basu09}. For most of the \\kepler\\ targets astrometric parallaxes will be available, while for \\corots\\ we must wait for the \\gaia\\ mission."
    },
    "1005/1005.0384_arXiv.txt": {
        "abstract": "We investigate the velocity dispersion of Pal~14, an outer Milky-Way globular cluster at Galactocentric distance of 71 kpc with a very low stellar density (central density 0.1-$0.2\\msun$/pc$^3$). Due to this low stellar density the binary population of Pal~14 is likely to be close to the primordial binary population.\\\\ Artificial clusters are generated with the observed properties of Pal~14 and the velocity dispersion within these clusters is measured as \\citet{Jordi09} have done with 17 observed stars of Pal~14. We discuss the effect of the binary population on these measurements and find that the small velocity dispersion of 0.38 km/s which has been found by \\citet{Jordi09} would imply a binary fraction of less than 0.1, even though from the stellar density of Pal~14 we would expect a binary fraction of more than 0.5. We also discuss the effect of mass segregation on the velocity dispersion as possible explanation for this discrepancy, but find that it would increase the velocity dispersion further.\\\\ Thus, either Pal~14 has a very unusual stellar population and its birth process was significantly different than we see in today's star forming regions, or the binary population is regular and we would have to correct the observed 0.38 km/s for binarity. In this case the true velocity dispersion of Pal~14 would be much smaller than this value and the cluster would have to be considered as ``kinematically frigid'', thereby possibly posing a challenge for Newtonian dynamics but in the opposite sense to MOND. ",
        "introduction": "Stellar populations with densities less than about $10^2$ stars/pc$^3$ have a binary fraction, \\begin{equation} f_{bin} = \\frac{N_{bin}}{N_{bin}+N_{sing}}, \\end{equation} of at least 50 per cent. For example, in the solar neighbourhood (density 1 star/pc$^3$) $f_{bin} \\approx 0.6$ \\citep{Duquennoy91} or in OB associations (density $< 0.1$ star/pc$^3$) $f_{bin} > 0.7$ \\citep{Kouwenhoven07}. In star-forming regions the binary fractions can be even higher, such as in the Taurus-Auriga groups (density about 10 stars/pc$^3$) where the fraction is near 100 per cent (e.g. \\citealt{Kroupa03}). But even in high-density star-forming regions, such as the 1 Myr old Orion Nebula cluster which has a density of $10^4$ stars/pc$^3$ \\citep{Hillenbrand98}, the binary fraction is near 50 per cent \\citep{Koehler06}. Stellar-dynamical models have demonstrated that the Orion Nebula cluster must be expanding and that consequently further destruction of binary systems has mostly ceased \\citep{Kroupa01b}. Only in globular clusters, which are dense ($>10^3$ stars/pc$^3$) and which have a long dynamical history over a Hubble time, are binary fractions observed to be less than 10-20 per cent \\citep{Hut92}. But this is probably understandable as a result of the destruction of binary systems, first in the very dense initial configuration which these objects typically formed with \\citep{Marks08, Marks10}, followed by long-term binary-star burning.\\\\\\\\ Pal~14 is an old globular cluster  with a mass of about 10$^4 \\msun$, a (three-dimensional) half-light radius of 34 pc, and is located at a distance of 71 kpc from the Galactic centre \\citep{Hilker06, Jordi09}. It has a low density (0.1-0.2 stars/pc$^3$) and so conformity with known stellar populations would imply it to have $f_{bin} \\approx 0.5$. Measurements of $f_{bin}$ would thus allow testing the dependence of the binary population on the physical conditions during star formation, as Pal~14 is a low-metallicity population that formed nearly a Hubble time ago \\citep{Jordi09}. Furthermore, \\citet{Jordi09} suggest that Pal~14 may be significantly mass segregated. How mass segregation can affect a velocity dispersion measurement has not been investigated, though.\\\\\\\\ However, Pal~14 is also interesting for testing gravitational theories. That binaries can affect the dynamical mass estimate of stellar populations has already been studied in the context of dwarf-spheroidal satellite galaxies by \\citet{Hargreaves96} and has recently been re-addressed by \\citet{Kouwenhoven08} and \\citet{Gieles10} for star clusters in the Milky Way disk. The measured (one-dimensional) velocity dispersion of Pal~14 is ($0.64\\pm0.15$) km/s or ($0.38\\pm0.12$) km/s, depending on whether the outlying ``star 15'' of a sample investigated by \\citet{Jordi09} is included or not, respectively. Now, if Pal~14 has a binary fraction of 50 per cent or more, then the velocity dispersion measured from a single spectroscopic snapshot may be significantly contaminated, i.e. inflated, because of the binary-star orbital motions of the tracer stars. A potential problem arising here is that the dynamical velocity dispersion of the cluster must be about 0.5 km/s for Newtonian dynamics to hold. So there is not much room for a binary-star contribution to the measured velocity dispersion. Alternatively, if the binary population is normal (i.e. canonical), then the true (binary-corrected) velocity dispersion of the virialised cluster may be much smaller than the Newtonian value. If this were the case, then a Newtonian crisis would emerge, such that Pal~14 would be ``kinematically frigid'', a situation which is not expected from current theoretical dynamics. A third possibility is given by the unreliability of low-number statistics in such systems, i.e. a sample of about 17 stars may simply be insufficient to determine the velocity dispersion of such a stellar system.\\\\\\\\ Thus, either Pal~14 is in virial equilibrium, such that its stellar population would be non-canonical (a binary fraction much lower than all other known populations). Or, the cluster has a canonical stellar population, but is then kinematically frigid, in violation of Newtonian dynamics. This would not be remedied by MOND, because according to MOND the cluster ought to have an even larger velocity dispersion than the Newtonian value \\citep{Baumgardt05}. The third explanation, of low-number statistics, would imply that velocity dispersion measurements of low-velocity stellar systems would have to be taken with much more caution than is done in practice nowadays.\\\\\\\\ With this contribution we assess the expected velocity dispersion of Pal~14 assuming various binary fractions and Newtonian dynamics to be valid (Section 2). These numerical experiments are compared to the observed velocities in the cluster in Section 3. The conclusions are contained in Section 4. ",
        "conclusions": "\\label{Sec:Conclusions} From our test of the observed velocity dispersion of the Milky-Way globular cluster Pal~14, we have seen that the binary fraction within Pal~14 has to be less than 0.2 in the case of star 15 not being considered a member, in order to be consistent with a velocity dispersion as low as (0.38$\\pm$0.12) km/s. With taking star 15 into account the maximum binary fraction consistent with the observational uncertainties is about 0.4. This poses a number of questions on the nature of Pal~14.\\\\\\\\ As a first explanation of these findings we may assume that the binary fraction within Pal~14 is indeed as low as found. But, as stated above, the density within Pal~14 is as low as 0.1-0.2 stars/pc$^3$. The effect of disruption of binaries due to dynamical stellar evolution within its age of about 11.5 Gyr is therefore negligible and we even may see here the primordial binary population of a globular cluster \\citep{Hasani10}. Thus, the formation of Pal~14 must have been significantly different from what is observed in all other star forming sites today. Or Pal~14 must have undergone a very dense and violent phase in which most of the binaries were burned, but as we can see in the ONC today, this is very unlikely to happen for a sufficiently long time span that the binary fraction drops below 0.5. Furthermore, such a scenario would give rise to the question how Pal~14 could have expanded that much, as with a half-light radius of about 34 pc it is one of the most expanded globular clusters of the Milky Way today. Recent numerical studies moreover show that this expansion is very unlikely to be of pure dynamical origin, since expansion takes place on a relaxation time scale and the relaxation time of Pal~14 is of order of a Hubble time \\citep{Hasani10}.\\\\\\\\ A second option would be that the binary population is normal, i.e. above 0.5, and that thus the observed velocity dispersion has to be corrected for the effect of binaries. As we have seen, this can be as much as a factor of 10 in the case of 17 stars drawn from a cluster with a binary fraction of 1.0. Since the lower mass limit of Pal~14 is determined to be about $7000\\msun$ \\citep{Jordi09} this would imply that the true velocity dispersion of Pal~14 is much lower than the Newtonian prediction, i.e. Pal~14 is ``kinematically frigid''. This would be inconsistent with our understanding of Newtonian gravity and could neither be explained by considering MOND to be valid in Pal~14.\\\\\\\\ Moreover, we found that mass segregation increases the observed velocity dispersion of a cluster even further, and thus cannot explain the low velocity dispersion of Pal~14. The observed, flattened mass function of Pal~14 on the other hand suggests that Pal~14 is significantly mass segregated \\citep{Jordi09}. Thus, its unsegregated, i.e. for the effect of mass segregation corrected, velocity dispersion may be even lower than the values reported by \\citet{Jordi09}, which would worsen the problem.\\\\\\\\ Thus, the present state of knowledge on Pal~14 is that either its binary fraction is highly abnormal, given its low density, or that it is significantly sub-virial. The former possibility would imply a non-standard star-formation event which formed Pal~14, while the latter indicates a problem understanding the dynamics of Pal~14.\\\\\\\\ In any case, this investigation has shown that one has to be cautious with low velocity dispersions derived from small samples of cluster stars. The observed value from such a sub-sample tends to be significantly larger than the true velocity dispersion. This effect gets larger with increasing size of the allowed velocity window about the mean radial velocity. This could be especially important in the outer parts of star clusters (e.g. \\citealt{Scarpa07}), as was also recently shown by \\citet{Kuepper10}.\\\\\\\\ More line-of-sight velocity measurements, e.g. as a 2nd epoch spectroscopic snapshot, may help to reduce the statistical uncertainties and improve the significance of findings on whether or not Newtonian dynamics is valid in Pal~14 (see also \\citealt{Gentile10})."
    },
    "1005/1005.5568_arXiv.txt": {
        "abstract": "We present skeleton studies of non-Gaussianity in the Cosmic Microwave Background temperature anisotropy observed in the five-year \\textit{Wilkinson Microwave Anisotropy Probe} (\\WMAP) data. The local skeleton is traced on the 2D sphere by cubic spline interpolation which leads to more accurate estimation of the intersection positions between the skeleton and the secondary pixels than conventional linear interpolation. We demonstrate that the skeleton-based estimator of non-Gaussianity of the local type ($f_{\\rm NL}^{\\rm local}$) - the departure of the length distribution from the corresponding Gaussian expectation - yields an unbiased and sufficiently converged likelihood function for $f_{\\rm NL}^{\\rm local}$.  We analyse the skeleton statistics in the \\WMAP\\ 5-year combined V- and W-band data outside the Galactic base-mask determined from the KQ75 sky-coverage. The results are consistent with Gaussian simulations of the the best-fitting cosmological model, but deviate from the previous results determined using the  \\WMAP\\ 1-year data. We show that it is unlikely that the improved skeleton tracing method, the omission of Q-band data, the modification of the foreground-template fitting method or the absence of 6 extended regions in the new mask contribute to such a deviation. However, the application of the Kp0 base-mask in data processing does improve the consistency with the \\WMAP1 results.  The $f_{\\rm NL}^{\\rm local}$-likelihood functions of the data are estimated at 9 different smoothing levels. It is unexpected that the best-fit values show positive correlation with the smoothing scales. Further investigation argues against a point-source or goodness-of-fit explanation but finds that about 30\\% of either Gaussian or \\fnl\\ samples having better goodness-of-fit than the \\WMAP\\ 5-year data show a similar correlation. We present the estimate $f_{\\rm NL}^{\\rm local}=47.3\\pm34.9$ ($1\\sigma$ error) determined from the first four smoothing angles and $f_{\\rm NL}^{\\rm local}=76.8\\pm43.1$ for the combination of all nine. The former result may be overestimated at the $0.21\\sigma$-level because of point sources. ",
        "introduction": "\\label{sec_intro}  Generic inflationary models predict that the initial conditions of the post-inflation universe can be described by a Gaussian random-phase field with nearly scale-invariant fluctuations. These subsequently seed the perturbations that generate Cosmic Microwave Background (CMB) anisotropies and structure formation thereafter. The Gaussianity of the statistics determined from measures of the CMB anisotropy and large scale structure distribution can provide evidence that validates the inflationary scenario of the extremely-early Universe. Besides the simplest single-scalar field model that predicts a truly Gaussian initial condition \\citep{Guth_1981, Bardeen_etal_1983, Mukhanov_etal_1992}, there are a number of inflationary models predicting non-Gaussianity in two broad classifications, the \\textit{equilateral} type and the \\textit{local} type. The detection of a specific type of non-Gaussianity can shed light on the fundamental physical properties of inflation.  In this paper, we are concerned with a \\textit{local} type non-Gaussianity of the ``simplest weak nonlinear coupling'' case \\citep{Komatsu_etal_2001} \\begin{equation} \\Phi(\\vecx) = \\Phi_{\\rm L}(\\vecx) + f_{\\rm NL}^{\\rm local}\\left[\\Phi_{\\rm L}^2(\\vecx)-\\langle\\Phi_{\\rm L}^2(\\vecx)\\rangle\\right], \\end{equation} where $\\Phi(\\vecx)$ denotes the primordial curvature perturbations and $\\Phi_{\\rm L}$ is its linear Gaussian part. The amplitude of the non-Gaussianity is parameterised by the dimensionless coupling constant $f_{\\rm NL}^{\\rm local}$ (\\fnl\\ hereafter). The first observational constraint on \\fnl\\ - $-3500 < f_{\\rm NL} < 2000$ at 95 \\% C.L. - % was discussed by \\citet{Komatsu_etal_2002} using the angular bispectrum computed from the four-year \\textit{COBE} DMR data \\citep{Bennett_etal_1996}. A reduced bispectrum technique, hereafter the KSW estimator, \\citep{Komatsu_etal_2005}, was applied to the first-year and three-year \\WMAP\\ data, leading to $-58 < f_{\\rm NL} < 134$ \\citep{Komatsu_etal_2003} and $-54 < f_{\\rm NL} < 114$ \\citep{Spergel_etal_2007}, respectively. \\citet{Yadav_Wandelt_2008} employed an apparently improved estimator \\citep{Yadav_etal_2008} to obtain $27<f_{\\rm NL}<147$ for the V+W-band data outside the Kp0 mask with $\\ell_{\\rm max}=750$ with the  three-year \\WMAP\\ data. The \\WMAP\\ team used the same estimator to measure \\fnl\\ from the five-year \\WMAP\\ V+W-band outside the KQ75 mask with $\\ell_{\\rm max}=700$ and obtained $-9<f_{\\rm NL}<111$.  The possibility of detecting CMB non-Gaussianity using a group of morphological statistics, - Minkowski functionals (MFs) \\citep{Matsubara_2003, Hikage_etal_2006} - has also been studied. The departure of MFs from their Gaussian expectations has been tested to be an unbiased estimator for \\fnl\\ and then applied to the \\WMAP 3-year Q+V+W combined map yielding $-70<f_{\\rm NL}<91$ at the 95\\% C.L. \\citep{Hikage_etal_2008}. The \\WMAP\\ team re-investigated the MFs estimator with the 5-year template-cleaned V+W map outside the KQ75 mask, yielding $f_{\\rm NL}=-57\\pm60$ (68\\% C.L.) at resolution $N_{\\rm side}=128$ and $f_{\\rm NL}=-68\\pm69$ at $N_{\\rm side}=64$ \\citep{Komatsu_etal_2009}. It is still unclear why the MFs favour a negative best-fit amplitude for \\fnl\\ while the bispectrum estimator prefers a positive one, even though  the MFs can be formed by the weighted sum of the bispectrum. Thus it is of great importance to use different estimators to identify and investigate the weak non-Gaussian signal in \\WMAP\\ observations. In fact, the one-point probability density function (1-pdf) of the smoothed temperature field can also be implemented \\citep{Bernardeau_etal_2002} as an alternative non-Gaussianity estimator \\citep{Jeong_etal_2007}. Indeed, as  noted by \\citet{Novikov_etal_2006}, the normalised differential length of the skeleton is closely linked to this quantity, but the skeleton remains of interest due to its different sensitivity to specific aspects of the data, eg. the noise distribution. It is likely that a complete understanding of the data can only be realised after the application of a wide range of statistical tests.  The skeleton has been considered as a probe of the filamentary structures of a 2D or 3D smooth random field. The original definition of the skeleton is non-local, making the analytical discussion difficult and the numerical evaluation costly. \\citet{Novikov_etal_2006} first proposed a local approximation that ``the local skeleton is given by the set of points where the gradient is aligned with the local curvature major axis and where the second component of the local curvature is negative''. They also presented a numerical approach to trace the local skeleton and found an approximate expression for the differential length distribution of a Gaussian field. As another morphological statistical test, the method has been applied to both large-scale structure measures \\citep{Sousbie_etal_2006, Sousbie_etal_2008} and CMB anisotropies \\citep{Eriksen_etal_2004}. The latter was performed on the Q+V+W map of the first-year \\WMAP\\ data outside a base-mask that is defined on the Kp0 sky-coverage. Comparing with Gaussian simulations, the length distribution of the skeleton did not show significant deviation from the Gaussian predictions. The impact of non-excluded point sources was found to be small for the statistics concerned.  In parallel to studies of non-Gaussian signal estimators, several algorithms of simulating non-Gaussian realisations have been developed. \\citet{Komatsu_etal_2003} first simulated the local-type non-Gaussian component by integrating the spherical harmonics of $\\Phi_{\\rm L}^2(\\vecx)-V_{x}^{-1}\\int\\,d^3\\vecx \\Phi_{\\rm L}^2(\\vecx)$ in spherical harmonic space. Another strategy has been developed in which a pre-computed `filter' encoding the correlation properties of Gaussian curvature perturbation multipoles boosts the computation of high-resolution temperature and polarisation Gaussian and corresponding non-Gaussian maps \\citep{Liguori_etal_2003, Liguori_etal_2007}. This method was recently improved by \\citet{Elsner_etal_2009}. Such \\fnl\\ simulation methods provide the community with powerful tools to investigate the primordial non-Gaussianity and the impact of other astrophysical and systematic effects on it.  In this paper, the skeleton length distribution is adopted as an estimator of the local-type non-Gaussianity. We adopt the cubic spline interpolation to trace the underlying local skeleton rather than the conventional linear one to make a more accurate estimation of the intersection position between the skeleton and pixel edge. Motivated by MFs studies on \\fnl, the statistical properties of the skeleton length distribution and the convergence of an \\fnl\\ estimation methodology are investigated from the \\fnl\\ simulations. We then analyse the skeleton statistics in the five-year release of the \\WMAP\\ data and compare with both Gaussian and non-zero \\fnl\\ samples. The results of the null Gaussian test are compared with those of \\citet{Eriksen_etal_2004} for the first-year \\WMAP\\ data, and then we use the skeleton estimator to compute a likelihood estimate for \\fnl.  This paper is organised as follows. In Section \\ref{sec_method}, we carry out numerical studies on the CMB local skeleton, including the skeleton statistics utilised in our analysis (Section \\ref{subsec_sta}) and the test of unbiasedness and convergency of \\fnl-likelihood led by skeleton estimator from noise-free \\fnl-simulations (Section \\ref{subsec_nulltest}). \\ref{subsec_wmap_data} presents an overview of the \\WMAP\\ data and the instrumental properties that should be encoded into our simulations to make an unbiased comparison and parameter estimation. Section \\ref{subsec_analysis} describes the process of computing the estimator and further analysis from both the observed data and simulations having  consistent instrumental properties and sky-coverage. Results are reported in Section \\ref{sec_results}, including the analysis and discussion of a Gaussian frequentist test (Section \\ref{subsec_gauss_res}) and \\fnl-estimations (\\ref{subsec_fnl_res}). Finally, we present our conclusions in Section \\ref{sec_conclusions}. ",
        "conclusions": "\\label{sec_conclusions}  In this paper, we have studied the local-approximation to the skeleton on a 2D sphere pixelised in the HEALPix scheme, and refined the method of tracing the quantity.  The statistical properties of the skeleton estimator have subsequently been investigated using mock CMB temperature anisotropy maps.  The cubic spline interpolation method locates the skeleton knots more accurately than the simple linear method, which makes the local linear system more robust at the knots. This is of great importance for finer analysis of the local system. For example, the studies on skeleton classification \\citep{Pogosyan_etal_2009}, which is performed by analysing the eigenvalues of the linear characteristic equation, request highly accurate estimation of such eigenvalues in particular around the demarcation point between two types of skeleton. Our modification provides a more reliable basis for this kind of study. The departure of the skeleton length distribution from its Gaussian expectation shows connections with both the sign and the magnitude of \\fnl\\ so that it would yield a \\fnl-likelihood function. Based on simulated sets of CMB temperature anisotropy with a local type  of non-Gaussian component, it has been tested that both the differential and cumulative skeleton estimators provide unbiased and sufficiently converged likelihood function for \\fnl, but the latter yields a likelihood 10\\% more converged than the former.  The estimator was applied to the five-year \\WMAP\\ data release and the results compared with not only the Gaussian predictions, but also the results from the first-year \\WMAP\\ data processing by \\citet{Eriksen_etal_2004}. An \\fnl-likelihood function has been estimated by computing the $\\chi^2$ on the basis of 2500 sets of \\fnl\\ samples. We have also investigated the goodness of fit, the impact of the point sources and the comic variance effect on the best-fit amplitudes of \\fnl. The analysis is carried out on the V+W combined map for various sky coverages.  The processing steps in our analysis follow closely those of \\citet{Eriksen_etal_2004} but utilise the new five-year KQ75 mask and combined V-  and W-band data. The smoothing scales adopted in our data processing are also identical to those selected in \\citet{Eriksen_etal_2004}. Our skeleton results show an apparent deviation from the first-year ones. According to an extensive series of tests, it is the difference between the two Galactic plane regions defined by the KQ75 and Kp0 masks that contributes mostly to the shifts. Generally, the KQ75 mask excludes a more extended region close to the Galactic plane than the Kp0 mask, and this should be more conservative for temperature analysis. This kind of deviation to the skeleton estimates implies a systematic bias in \\fnl-estimation, in other words, previous \\fnl\\ studies carried out on Kp0 sky coverage (or even the less conservative KP2 mask) may be biased by the residual Galactic foreground within the dark-grey regions as shown in Figure \\ref{fig_base_mask}. We do not exclude the pixels located in point sources to allow sufficient convergence of the likelihood. The impact of the point sources on the estimator is investigated by analysing the difference to samples using median-filtered maps. The results show that the point sources do have an asymmetric impact on the estimator between the positive and negative temperature thresholds on the four smallest smoothing scales. However, the effect is less significant for larger FWHMs. The results of our frequentist analysis show that the \\WMAP5 data are consistent with Gaussian predictions.  We have estimated the \\fnl\\ likelihoods on 9 smoothing levels. The results show an unexpected positive-correlation between the best-fit amplitudes, $f_{\\rm NL}^{\\rm best}$, and FWHM smoothing scales. The peak of the likelihood function seems to be `pulled' to a highly non-Gaussian region with the Gaussian case, $f_{\\rm NL}=0$, being `expelled' to the very tail of likelihood for some large smoothing scales. Further investigations argue against a point source explanation since the median-filtered data still exhibit such a correlation. However, the presence of point sources may yield an over-estimation of $f_{\\rm NL}^{\\rm best}$.  The combination of samples for the first 4 and all 9 smoothing scales lead to the best-fit amplitudes with $1\\sigma$ errors, $f_{\\rm NL}=47.3\\pm34.9$ and $f_{\\rm NL}=76.8\\pm43.1$, respectively. The median-filter studies suggest that the best-fit over 4 scales may be over-estimated at the $0.21\\sigma$-level because of point sources. An investigation has been carried out on the unexpected correlation between $f_{\\rm NL}^{\\rm best}$s and smoothing scales using both Gaussian and \\fnl\\ samples with a goodness-of-fit better than that for \\WMAP5. About 30\\% of them show the behaviour seen in our analysis, so that cosmic variance may be an appropriate explanation for this issue."
    },
    "1005/1005.2986_arXiv.txt": {
        "abstract": "A possible slowing down of the cosmic expansion is investigated through a cosmographic approach. By expanding the luminosity distance to fourth order and fitting the SN Ia data from the most recent compilations (Union, Constitution and Union 2), the marginal likelihood distributions for the deceleration parameter today suggest a recent reduction of the cosmic acceleration and indicate that there is a considerable probability for $q_0>0$. Also in contrast to the prediction of the $\\Lambda$CDM model, the cosmographic $q(z)$ reconstruction permits a cosmic expansion history where the cosmic acceleration could already have peaked and be presently slowing down, which would imply that the recent accelerated expansion of the Universe is a transient phenomenon. It is also shown that to describe a transient acceleration the luminosity distance needs to be expanded at least to fourth order. The present cosmographic results depend neither on the validity of general relativity nor on the matter-energy contents of the Universe. ",
        "introduction": "Independent measurements by various groups have suggested that the recent expansion of the Universe is speeding up and not slowing down, as believed for many decades \\cite{perlmutter_riess}. In other words, by virtue of some unknown mechanism, the cosmic expansion underwent a ``dynamic phase transition\" whose main effect was to change the sign of the universal deceleration parameter $q(z)$ at low redshifts. Together with other complementary observations, such results are strongly indicating that we live in a flat, accelerating Universe composed of $\\sim$ 1/3 of matter (baryonic + dark) and $\\sim$ 2/3 of an exotic component with large negative pressure, usually named dark energy (see \\cite{reviews} for some reviews).  However, what seems more certain is that the Universe was accelerating in a recent past, but that its current state of acceleration is less well determined \\cite{shapiro}. Some authors, using a dynamical ansatz for the dark energy equation of state, have suggested that the cosmic acceleration has already peaked and that we are currently witnessing its slowing down \\cite{2009PhRvD..80j1301S}. This behavior was also suggested by other works combining different methods and data \\cite{WuYu}.  The possibility of a transient cosmic acceleration appears in several theoretical scenarios \\cite{Others} and is also theoretically interesting because eternally accelerating universes (like $\\Lambda$CDM) are endowed with a cosmological event horizon which prevents the construction of a conventional S-matrix describing particle interactions. Indeed, such a difficulty has been pointed out as a severe theoretical problem for any eternally accelerating Universe \\cite{HKS01,CALR06}. Naturally, whether the acceleration of the Universe is slowing down the overall dynamic behavior in the future may be profoundly different from the one predicted by the $\\Lambda$CDM evolution.  Here we investigate the question related to a possible slowing down of the recent accelerating stage of the Universe by using a distinct method. In order to access the present value of the deceleration parameter we consider the so-called cosmographic approach \\cite{2004ApJ...614....1J,Visser04,Visser05,2007CQGra..24.5985C,2008PhRvD..78f3501C,2009arXiv0904.3550G,2010JCAP...03..005V,2010Ap&SS.330....7J} instead of assuming the validity of a specific gravitational theory together with a given dynamical dark energy model. As happens with kinematic models \\cite{shapiro,2009arXiv0904.3550G,TurRie02,riess04,EM1,EM,rapp07,Gong,CL08,Cunha09}, the basic advantage of the cosmographic approach over dynamic models is that the former depends neither on the validity of general relativity nor on the matter-energy contents of the Universe. In addition, to verify the dependence of the method on a particular set of supernova data, we repeat our analysis by considering three different samples, namely, the Union \\cite{Kowalski}, Constitution \\cite{2009ApJ...700.1097H}, and Union 2 \\cite{Union2} compilations. ",
        "conclusions": "To address the main question of this work, we should focus on the results for the Union 2 compilation, since it is currently the largest and most updated. Both kinds of cosmographic parametrization, $d_L^{(4)}(z<z_{cut})$ and $d_L^{(4)}(y)$, allow the circumvention of convergence issues related to the power expansion, broadly agree and (i) suggest that lower redshift SN events point to lower cosmic acceleration, and (ii) indicate that there is a considerable probability for the Universe to be decelerating today. The slowing down of the cosmic acceleration and even a transient acceleration scenario is allowed in all cosmographic parametrizations.  The possibility that those results may be due to bias induced by the form of the parametrizations and data analysis artifacts may be examined inside the cosmographic approach by going to higher order in the scale factor expansion and through the use of simulated data inside some conveniently chosen fiducial cosmological models. That in itself would constitute an interesting follow up study that could be useful to improve the understanding of the limitations of the observational data sets and of the analysis methods.  What may be causing a possible slowing down of cosmic acceleration and transient cannot be answered with basis only in a cosmographic approach. We recall that in the accepted scenario (inflation followed by radiation, dark matter and dark energy dominated stages), the Universe already had two accelerating stages and what is being discussed is whether the present supernova data are compatible with a new transition to a decelerating phase. Some examples of dynamical scenarios discussed in the literature include quintessence models in which the equation of state parameter becomes positive in the future, and some D-branes inspired scenarios, among others \\cite{Others,CALR06}. Naturally, we cannot exclude the possibility that such a slowing down might be provoked by some exotic kinematic (or dynamic effect) in the presence of inhomogeneities \\cite{Buchert}.  In conclusion, we used fourth order expansions of the luminosity distance to describe the recent cosmic expansion as probed by SN Ia observations. We showed that one needs to go at least to fourth order in the expansion of the luminosity distance to be able to describe a transient acceleration. Other works \\cite{2008PhRvD..78f3501C,2010JCAP...03..005V}, using smaller data sets than some of those used here, indicate that going to higher orders does not add meaningful statistical information. Under a cosmographic analysis of the most recent SN Ia compilations, we have found a suggestion of cosmic acceleration slowing down and the existence of a considerable probability in the relevant parameter space that the Universe is already in a decelerating expansion regime. This result is in great contrast to the standard $\\Lambda$CDM model, which predicts that the cosmic expansion is accelerating and must remain so forever with an ever increasing acceleration.  The predictability of the cosmographic parametrizations are considerably inferior to the $\\Lambda$CDM model and other dynamical models, but that seems to be the price for avoiding the usual assumptions about the physical contents of the Universe and the underlying gravity theory. It should also be stressed that the price for avoiding convergence issues for the expansion in powers of $z$ has severe consequences on model predictability, but this could be overcome by a much larger SN Ia sample (which is foreseen in future surveys), by the use of other kinds of data and priors, and maybe by a smarter (optimal) cosmographic modeling."
    },
    "1005/1005.3253_arXiv.txt": {
        "abstract": "We present {\\it Suzaku} observations of five hard X--ray selected  nearby Seyfert 2 galaxies. All the sources were clearly detected with the {\\tt pin} Hard X--ray Detector up to several tens of keV, allowing for a fairly good characterization of the broad--band X--ray continuum. We find that a unique model, even including multiple components, fails to represent the spectra of all the sources. Heavy obscuration manifests itself in different flavours. For two sources there is evidence for a reflection dominated continuum; among the other three, one is ``mildly\" Compton thick ($N_H \\sim$ 10$^{24}$ \\cm), while the remaining two are heavily obscured ($N_H \\simeq 10^{23.5}$ cm$^{-2}$), but Compton thin. Strong, narrow, iron K$\\alpha$ lines ($EW \\sim$ 1--2 keV) due to neutral or mildly ionized gas, are detected in Compton thick AGN. In all of them the K$\\alpha$ line is accompanied by the K$\\beta$. The intensity and shape of the soft X--ray spectrum are different from object to object. Soft X--rays may  originate from a nuclear component scattered off, or leaking through, the X--ray absorber, plus thermal X--rays from the host galaxy. Emission from circumnuclear gas photoionized by the active nucleus, parameterized with a power law plus individual narrow Gaussian lines, also provides an acceptable description of the soft X--ray spectra. The limited {\\it Suzaku} {\\tt XIS CCD} energy resolution does not allow us to draw firm conclusions on the origin of the soft X--ray emission. We briefly discuss our findings in the light of AGN Unified model and the geometry of the obscuring gas. ",
        "introduction": "A fraction as high as 50\\% of Seyfert 2 galaxies in the nearby Universe are obscured in the X--ray band by column densities of the order of, or larger than the inverse of the Thomson cross-section ($N_H\\ge \\sigma_T^{-1} \\simeq 1.5 \\times 10^{24}$~cm$^{-2}$), hence dubbed Compton thick (CT). If the optical depth ($\\tau = N_H \\sigma_T$) for Compton scattering does not exceed values of the order of ``{\\it a few}\", X--ray photons with energies higher than 10--15 keV are able to penetrate  the obscuring material and reach the observer. For higher values of $\\tau$, the entire X--ray spectrum is depressed by Compton down scattering and the X--ray photons are effectively trapped by the obscuring material irrespective of their energy. The former class of sources (mildly CT) can be efficiently detected by X--ray instruments sensitive above 10 keV, while for the latter (heavily CT) their nature may be inferred  through indirect arguments, such as the presence of a strong iron K$\\alpha$ line over a flat reflected continuum. The search for and the characterization of the physical properties of CT AGN is relevant to understand the evolution of accreting Supermassive Black Holes (SMBHs). In particular, mildly CT AGN are the most promising candidates to explain the so far largely unresolved spectrum of the X--ray background around its 30 keV peak (Worsley et al. 2005; Treister \\& Urry 2005; Gilli et al. 2007). According to the Gilli et al. (2007) XRB synthesis model, their integral contribution to the hard X--ray background is of the order of 25--30\\%. This fraction has been estimated under simplified hypotheses. In particular, the same luminosity function and cosmological evolution of unobscured and Compton thin AGN is assumed; moreover, the number density of ``heavily\" CT AGN is the same of ``mildly\" CT. While these assumptions are not inconsistent with the present observational framework (Risaliti et al. 1999; Guainazzi et al. 2005), it should be noted that absorption column densities  in excess of 10$^{24}$ \\cm\\ were measured or inferred for about a few tens of nearby AGN (Comastri 2004; Della Ceca et al. 2008), and only a handful of them are known beyond the local Universe (Norman et al. 2002; Iwasawa et al. 2005). CT AGN may represent a significant fraction of the accretion power in the Universe and indeed a correction for the CT contribution to the estimates of the SMBH local mass density has to be included in the calculations (Marconi et al. 2004).  An unbiased census of extremely obscured AGN would require to survey the hard X--ray sky above 10 keV with good sensitivity. Such an argument is one of the key scientific drivers of the {\\sc nustar} (Harrison et al. 2010) and {\\sc astro-h} (Takahashi et al. 2010) missions, which will be launched in the next few years, and of the {\\sc nhxm} mission study (Pareschi et al. 2010).  For the time being, one has to rely on the observations obtained by the high--energy detectors on board {\\sc bepposax}, {\\sc integral}, {\\sc swift} and, more recently, {\\it Suzaku}. All--sky surveys were performed using both the IBIS coded-mask telescope onboard {\\sc integral} and the BAT detector on board {\\sc swift}. Though limited to bright and thus low--redshift sources, they have proven to be quite successful. More than one hundred AGN are reported in both {\\sc integral} (Beckmann et al. 2009) and {\\sc swift}  catalogues (Tueller et al. 2010; Cusumano et al 2010).  Hard X--ray selection is less biased against absorption and thus provides a useful benchmark to study the column density distribution of obscured AGN and the fraction of CT sources in the local Universe. The large majority of {\\sc Integral} and {\\sc Swift} sources were observed by XMM and {\\it Chandra} and their  broad band X--ray spectra discussed by Winter et al. (2008). While the already known, nearby  CT AGN were recovered by {\\sc Swift} and {\\sc Integral} surveys, the fraction of ``newly\" discovered CT AGN is surprisingly low and apparently inconsistent, by about a factor of 2, with that predicted by Gilli et al. (2007) in the local Universe. On the basis of these findings, it has been proposed (Treister et al. 2009) that the contribution of CT AGN to the hard X--ray background may be significantly lower (by a factor of 2 to 3) than previously thought, with important implications for the evolution of the accretion power.  Since absorption column densities are often measured on relatively poor quality X--ray spectra taken from non-simultaneous observations, above and below about 10 keV, and combining different instruments, the CT fraction and, more in general, the absorption distribution in the local Universe may still be subject to several uncertainties. Good quality simultaneous spectra extending over the 0.5--100 keV energy range are needed for a robust measurement of absorption column densities, especially in the Compton thick regime. Moreover, nearby obscured AGN always show excess emission above the extrapolation of the obscured nuclear spectrum (e.g. Turner et al. 1997). Soft X--ray spectroscopy is a powerful tool to study the origin of this component which is plausibly related to warm gas photoionized by the nuclear continuum (Guainazzi \\& Bianchi 2007). The X--ray detectors onboard the Japanese satellite {\\it Suzaku} are well suited to this purpose, at least as far as ``local\" AGN are concerned. We have conceived a program with {\\it Suzaku} to observe nearby, relatively X--ray bright ($> 10^{-11}$ \\cgs) AGN selected by requiring {\\bf i)} a significant detection above 10 keV by either {\\sc Swift} or {\\sc Integral} and {\\bf ii)} evidence for significant X--ray absorption at lower energies from archival observations. The immediate science objective is the characterization of the broad band X--ray spectra of hard X--ray selected, heavily obscured AGN and a first step towards a better census of Compton thick absorption in the nearby Universe. ",
        "conclusions": "Thanks to the broad band energy range and sensitivity of the X--ray detectors onboard {\\it Suzaku} it has been possible to study, with good accuracy, the spectral properties of a sample of five hard X--ray selected type 2 AGN. The presence of cold obscuring gas has been unambiguosly established in all the sources. Although this result was expected, given the sample selection criteria, the X--ray spectra show a high degree of spectral complexity. In particular, the absorption column density and the relative intensity of the various spectral components differ from object to object.  The primary emission is seen, in three out of five sources, through column densities in the range $\\sim 10^{23.5-24}$ \\cm. The fraction of the intrinsic flux scattered into the line of sight and parameterized by $f_{scatt}$ in the partial covering fits is relatively low and the reflection component intensity weak. On the basis of {\\it Suzaku} observations of six {\\it Swift} selected AGN, thus an almost identical selection criterium,  Eguchi et al. (2009) found that sources with a low scattering fraction ($f_{scatt} <$  0.5 \\%) also have a relatively strong reflection component ($R$ $\\gsimeq$ 1). They dubbed these objects as ``new\" type AGN (see also Ueda et al. 2007) and proposed they may represent the tip of the iceberg of a hitherto uncovered population of AGN obscured by a geometrically and optically thick torus (Levenson et al. 2002). \\pn The analysis of the present data does not support the ``new\" type AGN interpretation. A possible explanation may be due to the strong degeneracy in the fit parameters. For example, if  the reflection component is also absorbed (like in the Eguchi et al. 2009), a solution in terms of a stronger reflection intensity and lower normalization of the scattered component is found, and results similar to Eguchi et al. (2009) are recovered also in our sample (Comastri et al. 2009). However, we believe that the quality of our data does not allow us to constrain absorption on the reflection component; moreover, our best fit solution has the advantage to have a lower number of free parameters. It is also worth pointing out that the intensity of the reflection spectrum from Compton thick gas strongly depends on the geometry of the system. The reflection fraction, obtained from fitting  disc--reflection models to Compton thick AGN, cannot be directly associated to a solid angle (Murphy \\& Yaqoob 2009). A better understanding of the nature of ``new\" type AGN and whether there are two distinct classes of Seyfert 2, possibly related to the geometry of the obscuring medium, or a continuous distribution of spectral parameters requires a larger sample of objects with a good counting statistic to allow for a detailed spectral analysis. \\pn The observational evidence of obscured AGN with a low scattering fraction is at variance with the average values adopted for the above parameters in the Gilli et al. (2007) XRB synthesis models. We have tested the impact of different normalizations for the scattering component in absorbed ($N_H > 10^{22}$ \\cm ) AGN. The fit to the XRB spectrum and source counts are basically unaffected, unless $f_{scatt}$ exceeds values of the order of 10 \\%. High $f_{scatt}$ values appears to be the exception, rather than the rule, and are definitely ruled out among highly obscured {\\it Swift} selected AGN. \\pn The possibility that photoionized gas is responsible for the soft X--ray emission in obscured AGN is consistent with the present observations. It is interesting to note that the residuals in the soft X--ray band of the NGC 5728 {\\it Chandra} spectrum (Zhang et al. 2006) and ESO 137--G34 XMM--{\\it Newton} spectrum (Malizia et al. 2009), both fitted with thermal emission, suggest the presence of unaccounted for emission lines at energies close to those reported here. Unfortunately, the physical status of the ionized gas cannot be constrained by the present observations. Only a limited number of emission lines are individually detected above the 3$\\sigma$ level. It is however reassuring that the best fit energies reported in Table~4 are consistent with those expected by ionised metals observed with the XMM--{\\it Newton} and {\\it Chandra} gratings observations of much brighter Seyfert 2 galaxies (Brinkman et al. 2002; Kinkhabwala et al. 2002; Guainazzi \\& Bianchi 2007). The $\\sim$ 0.9 keV feature, identified with {\\sc Ne ix} appears to be ubiquitous, possibly with the exception of NGC 4992 which is, however, the faintest soft X--ray source in the sample.  It is worth noting that the three brightest soft X--ray sources in the sample, all of them Compton thick (Fig.~1), show strong {\\sc [oiii]} ionization cones in optical images (e.g. Wilson et al. 1993). Moreover, extended soft X--ray emission almost spatially coincident with {\\sc [oiii]} cones is detected by {\\it Chandra} imaging observations of NGC 5728 (Zheng et al. 2006). Furthermore, ESO 137--G34 and ESO 323--G32 are hosted by S0 galaxies which are expected to have neither intense star formation activity, nor thermal X--ray emission from hot gas.  Although it may be premature to discard, on the basis of the present X--ray observations, a thermal origin for the soft X--ray emission, we stress that the above described arguments would support the presence of photoionized plasma, which is also favored by high resolution {\\sc rgs} observations of bright obscured AGN (Guainazzi \\& Bianchi 2007). \\pn Long--look observations of selected sample of obscured AGN are clearly needed to elucidate the nature of the soft X--ray emission in these objects, though firm conclusions can be drawn only by high resolution (with either gratings or calorimeters) X--ray observations."
    },
    "1005/1005.4473_arXiv.txt": {
        "abstract": " ",
        "introduction": "With the development of observation technology, we have more opportunity to observe and analyze the thermal radiation from the compact star surfaces. The comparison between the predictions of theoretical study of the thermal evolution  and observations of thermal emission from the compact stars gives a potentially powerful method to study fundamental properties of superdense matter in compact star interiors. \\cite{hpy07, lp07}.  It is generally believed that dense matter can be superfluid at such high density and low temperature \\cite{ls01}. Baryon pairing (or quark pairing) strongly suppresses not only the neutrino emission but also the heat capacity of nuclear matter. Meanwhile, Cooper pair formation and breaking also will effect the neutron star cooling. Furthermore, as a neutron star spins down, some heating mechanisms may be present and will play important roles in the thermal evolution of NS. They have been extensively discussed, for example, rotochemical heating \\cite{rei95,fr05}, mutual friction between superfluid and normal components of the star \\cite{SL89}, crust cracking \\cite{cczc92} and the dissipation of rotational energy due to viscous damping \\cite{zheng2006}. We will focus on rotochemical heating in the following study.   It's well-known \\cite{ykgh01} that for nonsuperfluid NS with the core composed of standard nuclear matter, we have two distinct cooling regimes, slow and fast cooling. The transition from slow to fast cooling with increasing M occurs in a very narrow mass range and is sharp. After considering the superfluidity, the Urca neutrino emission and the NS heat capacity will be strongly suppresses, and neutron emission due to Cooper pairing will occur \\cite{frs76}. Then these effects will make NS warmer. Recently, when including nucleon superfluidity and using APR EOS , \\cite{gkyg05} found a particular scenario of neutron star cooling that the theoretical cooling models of isolated middle-aged neutron stars can be divided into three distinct types: slow, moderate and fast cooling. They neglected the heating effects, while we will take this effects into account which may make the cooling scenario different. ",
        "conclusions": ""
    },
    "1005/1005.4968_arXiv.txt": {
        "abstract": "We present the    clustering measurement of hard X-ray selected AGN in the local Universe. We used a sample of  199 sources  spectroscopically confirmed detected by {\\em Swift}-BAT in its 15-55 keV all-sky survey. We measured the real space projected auto-correlation function and detected a signal significant on projected scales lower than 200 Mpc/h. We measured a correlation length of r$_0$=5.56$^{+0.49}_{-0.43}$ Mpc/h and a slope $\\gamma$=1.64$_{-0.07}^{-0.08}$. We also measured the auto-correlation function of Type I and Type II AGN and found higher correlation length for  Type I AGN. We have a marginal  evidence of luminosity dependent clustering of AGN, as we detected a larger correlation length of luminous  AGN than that of low luminosity sources. The corresponding typical host  DM halo masses  of {\\em Swift}-BAT are $\\sim$ log(M$_{DMH}$)$\\sim$ 12-14 h$^{-1}$M/M$_{\\odot}$, depending on the subsample. For the whole sample we measured    log(M$_{DMH}$)$\\sim$ 13.15 h$^{-1}$M/M$_{\\odot}$ which is the typical mass of a galaxy group. We estimated that the local AGN population has a typical lifetime $\\tau_{AGN}\\sim$0.7 Gyr, it is powered  by SMBH with mass  M$_{BH}\\sim$1-10$\\times$10$^{8}$ M$_{\\odot}$  and accreting with very low efficiency,  log($\\epsilon$)$\\sim$-2.0. We also conclude that local AGN host galaxies are typically red-massive  galaxies with stellar mass of the order 2-80$\\times$10$^{10}$ h$^{-1}$M$_{\\odot}$. We  compared our  results with clustering predictions of  merger-driven AGN triggering   models and found a  good agreement. ",
        "introduction": "It is now  commonly believed that almost all galaxies host a central supermassive Black Hole (SMBH). Dynamical evidence show that the mass of the central BHs are closely linked to the mass as well as the stellar velocity dispersion of the bulge component of the host galaxy (Kormendy \\& Richstone 1995, Magorrian et al. 1998). This  suggests that the formation and evolution of the spheroidal component of galaxies and their SMBH are closely connected. It is of  upmost importance to understand the mechanism of funneling  interstellar gas into the  vicinities of the SMBH, triggering the accretion. Galaxy mergers or tidal interaction between close pairs may have played a major role (Hopkins et al. 2007), furthermore some mechanism internal to the galaxy, like galaxy disk instability may be important. Clustering properties of AGN in various redshifts give an important clue to understand which of these mechanisms trigger AGN activities in what stage of the evolution of the universe, through, e.g., the mass of the Dark Matter Halos (DMH) in which they reside which is linked to BH mass life time and Eddington rate. Measurements of the AGN clustering show that  AGN are typically hosted by DMH with a mass of the order of  log(M)$\\sim$12.0-13.5 M/M$_\\odot$ % (Yang et al. 2006;  Miyaji et al. 2007; Gilli et al. 2009; Coil et al. 2009; Hickox et al. 2009; Krumpe et al. 2010). However, these measurements have been produced by using  AGN samples obtained by optical and soft X-ray  (i.e. 0.5-10 keV) surveys. Optical and soft X-ray selections miss a major part of the SMBH accretion. In the optical band the selection of AGN is biased by galaxy starlight dilution and by dust absorption. Although luminous soft X-ray emission is  a signature  of the presence of an AGN,  absorbed sources can  be missed with a soft X-ray selection, either because they are intrinsically less luminous % (Hasinger et al. 2008) or because of the high column density. However, X-ray emission from these sources leaks out at higher energies (i.e $>$ 5-10 keV) where the efficiency  of instruments mounting X-ray focusing optics is low. For this reason hard X-ray  selected samples could provide clean and unbiased samples of AGN. The {\\em Swift}-BAT all-sky survey provides a spectroscopically complete (100 \\%) sample of local AGN detected in the 15-55 keV energy  band, with an unprecedented depth and characterization of the source properties, from redshifts to column densities. In this letter we present  the first study of  clustering of hard X-ray selected  AGN in the local Universe. Throughout this paper we will assume a $\\Lambda$-CDM cosmology with $\\Omega_m$=0.3, $\\Omega_{\\Lambda}$=0.7, H$_0$=100$h^{-1}$ km s$^{-1}$ Mpc and $\\sigma_8$=0.8. Unless otherwise stated errors are quoted at the 1$\\sigma$ level. ",
        "conclusions": "In this letter we report on the measurement of clustering of 199 AGN in the local Universe using the Swift/BAT all-sky survey sample. This result gives, for the first time, an unbiased picture of the z=0 DMH-galaxy-AGN coexistence/evolution. We  obtained  a correlation length r$_0$=5.56$^{+0.49}_{-0.43}$  Mpc/h  and $\\gamma$=1.64$_{-0.08}^{+0.07}$. We  measured the ACF for Type I and Type II AGN and found a significant difference in their correlation lengths.  We have measured a marginally significant higher r$_0$ for high luminosity AGN than for the low luminosity ones. We propose that the observed difference in Type I vs. Type II clustering is driven by the intrinsic higher $<L_X>$ of Type I AGN as we show a marginal evidence of a correlation between r$_0$ and $L_X$. We estimated the typical mass of the DMH hosting an AGN of the order  log(M$_{MDH}$)$\\sim$13.15 h$^{-1}$M/M$_{\\odot}$. In Fig.  \\ref{fig:bias} we show  the  bias-redshift  plane results from AGN and galaxy surveys (references in the figure). In the same plot we show the expected evolution of different DMH masses. We compared only bias values of studies that rely on the real space correlation function $\\xi(r)$ (values of $\\sigma_{\\rm 8, AGN/GAL}$  from Krumpe et al.  2010). This approach allows us to compare all different clustering studies on a common basis.  The majority of the X-ray  surveys agree with a picture where AGN are typically hosted in DM halos with mass of the order of 12.5 h$^{-1}$M/M$_{\\odot}<$ log(M$_{MDH}$)$<$13.5  h$^{-1}$M/M$_{\\odot}$ which is the mass of moderately poor group. Optically selected AGN instead reside in lower density environment and of the order of the  log(M$_{MDH}$)$\\sim$12.5  h$^{-1}$M/M$_{\\odot}$. Another interesting aspect is that   X-ray selected AGN samples (including ours)  cluster similarly to red galaxies and that LL AGN or type II AGN are found typically in less massive environments. On the contrary HL AGN and Type I AGN are hosted in  massive galaxies in massive DM halos (clusters).  We estimated that Swift-BAT AGN are powered by black holes with a typical ~mass  log(M$_{BH}/M_{\\odot}$) $\\sim$8.5, accreting at very low Eddington ratio (i.e. $\\sim$0.01 L$_{Edd}$) and that they are hosted by massive galaxies with mass of the order $\\sim2\\times10^{11}$ M/M$_{\\odot}$. These properties, except $\\epsilon$, scale with  $<L_X>$ and Type (HL and Type I are hosted in higher mass DMH in  more massive galaxies with bigger black holes). The upper limits on the duty cycle suggest that these AGN are shining since at least for 0.2-1.2 Gyr \\footnote{HL and type I estimate may be wrong because of the relatively young age of 10$^{14}$  (M/$M_{\\odot}$) DMH}.      We then tested the AGN merger-driven triggering paradigm by comparing the theoretical predictions for AGN clustering of the model of  Marulli et al. (2008) and Bonoli et al. (2009) with our results on the whole sample. The theoretical model is based on the assumption that AGN activity is triggered by galaxy mergers and the lightcurve associated to each accretion event is described by an Eddington-limited phase followed by a quiescent phase modeled after Hopkins et al. (2005).  Applied to the Millennium Simulation (Springel et al. 2005), such model has been shown to be successful in reproducing the main properties of the black hole and AGN populations (Marulli et al. 2008) and the clustering of optical quasars (Bonoli et al. 2009).  Using this model, we computed the expected correlation length of a sample of  simulated  AGN at z$\\sim$0 with intrinsic $L_{bol}$ luminosities similar to the ones of our observed AGN. The model predicts a clustering  length  r$_0$=5.68$\\pm{0.08}$ which is in agreement  within 1$\\sigma$ with our measurement.  By merging the observational evidences and the model predictions, a plausible scenario   for the history of local AGN is  the following: \\begin{itemize} \\item {\\em Swift-BAT} AGN switched on about 0.7 Gyr Ago after a galaxy merger event. \\item  They shine in an Eddington-limited regime for the first part of their lives where they gain most of their mass. \\item  In the second phase of their lives (i.e. after 0.2-0.5 Gyr) they start to accrete with lower and lower efficiency. Their luminosity drops because of the decreased gas reservoirs. \\item At z$\\sim$0 they have grown to $\\sim$10$^{8-9}$ M$/M_{\\odot}$ SMBHs, shining as moderately low-luminosity  AGN at low accretion rates. \\end{itemize}"
    },
    "1005/1005.3123_arXiv.txt": {
        "abstract": "{} {We present clumps of dust emission from Herschel observations of the Large Magellanic Cloud (LMC) and their physical and statistical properties. We catalog cloud features seen in the dust emission from Herschel observations of the LMC, the Magellanic type irregular galaxy closest to the Milky Way, and compare these features with \\hi\\ catalogs from the ATCA+Parkes \\hi\\ survey.} {Using an automated cloud-finding algorithm, we identify clouds and clumps of dust emission and examine the cumulative mass distribution of the detected dust clouds. The mass of cold dust is determined from physical parameters that we derive by performing spectral energy distribution fits to 250, 350, and 500 $\\mu$m emission from SPIRE observations using dust grain size distributions for graphite/silicate in low-metallicity extragalactic environments.} {The dust cloud mass spectrum follows a power law distribution with an exponent of $\\gamma=-1.8$ for clumps larger than 4$\\times$10$^2$ M$_\\odot$ and is similar to the \\hi\\ mass distribution. This is expected from the theory of ISM structure in the vicinity of star formation.} {} ",
        "introduction": "The overall structure and evolution of the ISM of a galaxy is determined by the physical and dynamical conditions in the diffuse ISM where most of the processing of gas and dust clouds occurs. Thus inspecting the cloud properties by studying the atomic gas and the dust emission under different environmental conditions will help us to understand how this processing of matter proceeds.  In the present paper, we study the dynamical conditions of the diffuse gas in the Large Magellanic Cloud (LMC). We examine optically thin dust emission at far-infrared (FIR) and sub-millimeter wavelengths by studying the spectral energy distributions (SEDs) inferred from radiative transfer models. By combining an \\hi\\ synthesis survey and an imaging survey in the FIR and submillimeter emission performed by {\\it Herschel} Space Observatory (Pilbratt et al. 2010), we investigate the detailed relationship between the diffuse atomic gas and interstellar dust emission in the LMC. The {\\it Herschel} Space Observatory is designed to study the submillimeter and FIR emission from our universe and offers unprecedented resolution. The {\\it Herschel} has a payload of two cameras, PACS (Poglitsch et al. 2010) and SPIRE (Griffin et al. 2010). PACS employs two bolometer arrays for imaging, and SPIRE comprises spider-web bolometers with NTD Ge temperature sensors. Mapping strategy and data processing procedures are described in detail by Meixner et al. (2010).  The {\\it Herschel} survey of the Magellanic Clouds will provide information about the interstellar dust emission (Meixner et al. 2006) from all phases of the ISM in the Magellanic Clouds. For diffuse atomic gas, we use data from the \\hi\\ aperture synthesis survey in Kim et al. (1998) and data from the combined survey using single-dish observations with the Parkes 64-m radio telescope (Staveley-Smith et al. 2003) that was published in Kim et al. (2003). The SED models provide spectral information in the 160 $\\mu$m to 500 $\\mu$m wavelength range. Relatively optically thin dust emission at FIR and submillimeter wavelengths can also provide information about the gas mass in the ISM because the gas and dust are well mixed in most of the ISM phases. Therefore, the dust emission at the FIR and submillimeter wavelengths can trace the column density and the structure of the ISM.  \\begin{figure} \\scalebox{0.45} {\\includegraphics{14645f1.ps}} \\caption[width=12cm]{Dust clouds seen in the SPIRE 500 $\\mu$m image were identified using an automatic finding algorithm in which the lowest threshold was set to be 3 $\\sigma$ (blue contours), and they are overlaid on the \\hi\\ aperture synthesis image (Kim et al. 2003) of the LMC. 1$\\sigma$ is set to 0.3 MJy/sr at 500 $\\mu$m. The black box indicates the observed region in the Science Demonstration Program of Herschel Space Observatory (Meixner et al. 2010).\\\\} \\end{figure}  \\begin{figure} \\scalebox{0.46} {\\includegraphics{14645f2}} \\caption[width=12cm]{Dust clouds seen in the SPIRE 350 $\\mu$m image were identified using an automatic finding algorithm, the lowest threshold being set to 3 $\\sigma$ (red contours), and overlaid on the H$\\alpha$ image of the LMC (Kim et al. 1999). The black box indicates the observed region in the Science Demonstration Program of Herschel Space Observatory (Meixner et al. 2010).\\\\} \\end{figure} ",
        "conclusions": "Dust clumps have been identified and cataloged in the {\\it Herschel} SPIRE survey of the LMC using an automated cloud-finding algorithm. The distribution of cold dust clumps is remarkably similar to the \\hi\\ clump mass distribution, sharing an index of mass distribution, $\\gamma=-1.8$. However, the dust clump mass spectrum in the lower mass regime follows a flatter power law than the Salpeter stellar IMF."
    },
    "1005/1005.2125_arXiv.txt": {
        "abstract": "{Protostellar jets are known to emit in a wide range of bands, from radio to IR to optical bands, and to date also about ten X-ray emitting jets have been detected, with a rate of discovery of about one per year.} {We aim at investigating the mechanism leading to the X-ray emission detected in protostellar jets and, in particular, at constraining the physical parameters that describe the jet/ambient interaction by comparing our model predictions with observations available in the literature.} {We perform 2D axisymmetric hydrodynamic simulations of the interaction between a supersonic jet and the ambient medium. The jet is described as a train of plasma blobs randomly ejected by the stellar source along the jet axis. We explore the parameter space by varying the ejection rate, the initial Mach number of the jet, and the initial density contrast between the ambient medium and the jet. We synthesized from the model the X-ray emission as it would be observed with the current X-ray telescopes.} {The mutual interactions among the ejected blobs and of the blobs with the ambient medium lead to complex X-ray emitting structures within the jet. The X-ray sources consist of several components: irregular chains of knots; isolated knots with measurable proper motion; apparently stationary knots; reverse shocks. The predicted X-ray luminosity strongly depends on the ejection rate and on the initial density contrast between the ambient medium and the jet, with a weaker dependence on the jet Mach number.} {Our model represents the first attempt to describe the X-ray properties of all the X-ray emitting protostellar jets discovered so far. The comparison between our model predictions and the observations can provide a useful diagnostic tool necessary for a proper interpretation of the observations. In particular, we suggest that the observable quantities derived from the spectral analysis of X-ray observations can be used to constrain the ejection rate, a parameter explored in our model that is not measurable by current observations in all wavelength bands.} ",
        "introduction": "\\label{Introduction}  In the past decade, high-energy emission from protostellar jets (originally suggested by \\citealt{pm81}) has been discovered, taking advantage of both the high spatial resolution of Chandra and the high effective area of XMM-Newton. To date, about ten X-ray emitting Herbig-Haro (HH) objects are known and the main properties of their emission are reviewed in Table \\ref{tab-HH-X} (see also \\citealt{bop07}).  \\begin{table*}[!t] \\begin{center} \\caption{Physical properties observed in X-ray emitting HH jets.} \\label{tab-HH-X} \\par \\begin{tabular}{lccccccccl} \\hline \\hline object   &  $t_{\\rm exp}$ & $cnts$       & $L_{\\rm X} $            & $T_{X}$   & $L_{\\rm j}$ & $z_{\\rm X}$ & $D$   & $LM/HM$ & Reference \\\\ &  [ks]          &              &[$10^{29}$ erg s$^{-1}$] & [$10^{6}$ K] & [arcsec]    & [AU]        &  [pc]  &         \\\\ \\hline HH 2     & $21.4$         & $11$         & $5.2$   \t           & $1$\t   & $2$     & $56000^c$   & $480$  &  LM & \\citet{pfg01} \\\\ HH 154   & $97$           & $63$         & $3$\t     \t           & $4$\t   & $5$     & $70-140$    & $140$  &  LM & \\citet{fbm06}$^d$ \\\\ HH 80/81 & $37.3$         & $46/63$      & $450/430$\t           & $1.5$\t   & $-$     & $515000$    & $1700$ &  HM & \\citet{ptm04} \\\\ HH 168   & $78$           & $-$          & $30$\t\t           & $6.5$\t   & $-$     & $-$\t   & $730$  &  HM & \\citet{ptu09}  \\\\ HH 210   & $838$          & $31$         & $10$\t                   & $0.8$         & $-$     & $37000^c$   & $450$  &  HM & \\citet{gfg06} \\\\ HH 540   & $795.8$        & $\\approx200$ & $0.4$                   & $6.6$\t   & $-$     & $<90$       & $450$  &  LM & \\citet{kfg05}$^a$ \\\\ HH 216   & $78$           & $8$          & $10$    \t           & $-$\t   & $4$     & $-$\t   & $2000$ &  -  & \\citet{lgm07} \\\\ DG Tau   & $90$           & $18; 9$$^b$  & $0.12$   \t           & $3.4$\t   & $5$     & $30$        & $140$  &  LM & \\citet{gsa08} \\\\ Z CMa    & $39.6$         & $20$         & $>2.5$   \t           & $2.3$\t   & $-$     & $> 2000$    & $1050$ &  HM & \\citet{sho09} \\\\ TKH 8    & $89.2$         & $28$         & $20$   \t           & $35.4$        & $-$     & $450-900$   & $450$  &  LM & \\citet{tkk04}$^e$ \\\\ \\hline \\end{tabular} \\end{center} \\smallskip $t_{\\rm exp}$ is the exposure time of the observations (or the total exposure time adding together several observations, as in the case of DG Tau); $cnts$ are the collected photons; $L_{\\rm X}$ is the X-ray luminosity in the $[0.3-10]$ keV band; $T_{X}$ is the best fit temperature; $L_{\\rm j}$ is the linear size of the X-ray source; $z_{\\rm X}$ is the distance of the X-ray source from the central protostar; $D$ is the distance of the HH object; $LM/HM$ indicates the low-mass or high-mass young stellar object from which the jet originates.  $^a$ Improved estimates of $L_{\\rm X}$ and $T_{X}$ from Joel Kastner and Ettore Flaccomio, private communication $^b$ The first value is associated with the South-Western jet and the second with the North-Eastern jet in DG Tau $^c$ Values derived from SIMBAD $^d$ See also \\citet{bfr03} $^e$ See also \\citet{tkh01}  \\end{table*}  The X-ray sources detected in HH jets are characterized by different morphologies, luminosities, and locations within the jet. In some cases, the X-ray emitting region is located at the base of the (optical) jet, near the protostar from which the jet originates. This is the case, for instance, for the low-mass (LM) HH 154 and DG Tau jets, both in Taurus (see \\citealt{bfr03, gsb05}). In other cases, the X-ray source is located further away from the protostar as, for example, for the high-mass (HM) HH 80/81 jet. As for the morphology, some of the X-ray sources cannot be resolved by current instruments (and appear point-like), some others appear to be elongated.  In one case (HH 154) a knotty X-ray source is found (resembling the knotty morphology commonly observed in the optical emission from HH jets) which consists of an elongated tail with a measurable proper motion away from the stellar source and an apparently stationary point-like source (\\citealt{fbm06}).  Current models, aimed at explaining the origin of X-ray emission from protostellar jets, considered a continuous supersonic jet propagating through a homogeneous interstellar medium (\\citealt{bop04, bop07}). These models succeeded in predicting X-ray sources associated with jets with luminosity and proper motion in good agreement with the observations. They failed, however, in reproducing the complex knotty morphology of the X-ray source detected in HH 154 (\\citealt{fbm06}).  In a previous paper (\\citealt{bop10}; hereafter Paper I), we proposed an improvement to these models by considering the scenario based on a pulsed jet, i.e. a jet characterized by an ejection velocity varying randomly in time, and interacting with an initially homogeneous ambient medium. The aim was to investigate the origin of the irregular knotty structure observed in protostellar jets in different wavelength bands and the complex interactions occurring among blobs of plasma ejected from the stellar source. Our analysis showed that the mutual interactions of blobs ejected at different epochs and with different speed naturally reproduce the irregular pattern of knots observed along the jet axis in many HH objects and lead to a variety of plasma components which cannot be described by models of a jet ejected with a sinusoidal variable velocity (e.g. \\citealt{rdk07}).  Here we further explore the model of a randomly ejected pulsed jet presented in Paper I with the aim to investigate the properties of its X-ray emission and how these properties may depend on the physical conditions in which the jet evolves. In fact, X-ray emission from jets is observed in HH objects at different evolutionary stages and with different masses: for instance, HH 154 originates from a LM binary system of class 0 and class I sources; DG Tau is a more evolved LM classical T Tauri star; HH 80/81 originates from a HM protostar. Given the wide range of conditions in which X-ray emission from jets originates, it is important to address the following issues: where and when is the X-ray emission expected to arise from HH jets? How common is the high-energy emission from HH jets?  The paper is organized as follows: in Sect. \\ref{The model} we present our model; in Sect. \\ref{Results} we show our results concerning the morphology and luminosity in the X-ray band; in Sect. \\ref{Discussion} we discuss our results and in Sect. \\ref{Conclusions} we draw our conclusions. ",
        "conclusions": "  \\begin{itemize}  \\item[$\\diamond\\,\\,$]{In all the cases, the interactions of the ejected plasma blobs with the surrounding medium produce X-ray emitting features. The main components contributing to the X-ray emission are: irregular chains of knots; isolated high speed knots; steady knots; reverse shocks; oblique shocks.}  \\item[$\\diamond\\,\\,$]{Light jets produce significant X-ray emission consistent with the levels observed, whereas heavy jets are characterized by very faint or no X-ray emission and emit mostly in the optical band.}  \\item[$\\diamond\\,\\,$]{In the case of light jets (leading to detectable X-ray emission), the X-ray luminosity is mainly determined by the ejection rate of plasma blobs, rather than by the jet Mach number: higher ejection rates are related to more energetic objects, thus leading to higher X-ray luminosities.}  \\item[$\\diamond\\,\\,$]{Our model predicts X-ray luminosity and best-fit temperature in nice agreement with most of observed X-ray emitting jets. In particular the HH 154 jet is well described as a pulsed jet with ejection rate $\\Delta t = 2$ yr. Our model represents the first attempt to describe all the X-ray emitting jets detected so far.}  \\item[$\\diamond\\,\\,$]{We found that most of the X-ray emission is located at the base of the optical jet where the plasma blob collisions are the most energetic and where, therefore, the probability to detect X-ray emission is the highest. This result is consistent with the evidence that almost all the LM objects show emission localized preferentially at the base of the jet.}  \\end{itemize}  In conclusion, our model explains why only a small fraction of HH jets has been detected in X-rays. Indeed detectable X-ray emission may arise only under favorite conditions, namely high ejection rates of plasma blobs and high density contrast between the perturbed ambient medium and the ejected blob.  We stress that erroneous interpretations of observations can easily be drawn if the X-ray sources (reverse shocks, stationary knots, etc.) are not monitored frequently, since they show fast variability over time-scales of a few years. On this respect, numerical models can provide strong diagnostic tools to interpret observations. To this end, we provided here model predictions on the morphology and characteristics of the X-ray source associated with HH 154 that may be revealed in future observations."
    },
    "1005/1005.2885.txt": {
        "abstract": "A possible gluon-condensate-induced modified-gravity model with $f(R) \\propto |R|^{1/2}$ has been suggested previously. Here, a simplified version is presented using the constant flat-spacetime equilibrium value of the QCD gluon condensate and a single pressureless matter component (cold dark matter, CDM). The resulting dynamical equations of a spatially-flat and homogeneous Robertson--Walker universe are solved numerically. This simple empirical model allows, in fact, for a careful treatment of the boundary conditions and does not require a further scaling analysis as the original model did. Reliable predictions are obtained for several observable quantities of the homogeneous model universe. In addition, the estimator $E_{G}$, proposed by Zhang et al. to search for deviations from standard Einstein gravity, is calculated for linear sub-horizon matter-density perturbations. The QCD-scale modified-gravity prediction for $E_{G}(z)$ differs from that of the $\\Lambda\\text{CDM}$ model by about $\\pm 10\\%$ depending on the redshift $z$. ",
        "introduction": "\\label{sec:introduction}  It has been suggested that the gluon condensate~\\citep{ShifmanVainshteinZakharov1979,Narison1996} of quantum chromodynamics [QCD] in the $q$--theory framework~\\citep{KV2008-statics,KV2008-dynamics,KV2010-CCP1} may lead to a particular type of $f(R)$ modified-gravity model~\\citep{KV2009-gluoncond}. The cosmology of this model was explored in \\citet{K2010-gluoncondcosm}. Here, the model is simplified even further, allowing for a detailed study of the corresponding cosmological equations and eliminating all arbitrariness.   The model of this article is, however, considered from a different point of view than the one used in the original article~\\citep{KV2009-gluoncond}. Here, the aim is not to derive the particular model used from the underlying theory (general relativity and QCD), but to consider the model \\emph{per se} and to test the hypothesis whether or not this model gives a reasonably accurate description of the observed ``accelerating universe'' \\citep[cf.][and references therein]{Riess-etal1998,Perlmutter-etal1999, Komatsu-etal2009}. In other words, the approach is in the sprit of Kepler rather than Newton.  Still, the model is not chosen at random from all possible $f(R)$ modified-gravity models, but obeys certain physically motivated selection criteria, as discussed in Section~\\ref{subsec:Selection-criteria}. The selected QCD-scale modified-gravity model is detailed in Section~\\ref{subsec:Representative-model}, with a few technical points relegated to Section~\\ref{subsec:Additional-remarks}.  The corresponding cosmological model with a single matter component (cold dark matter, CDM) is introduced in Section~\\ref{subsec:From-three-components-to-one}. The resulting differential equations are given in Section~\\ref{subsec:Dimensionless-equations}. Three observable quantities of the homogeneous cosmological model are then discussed in Section~\\ref{subsec:Observables-quantities}. A fourth observable, this time of the inhomogeneous cosmological model, is also considered, namely, the estimator $E_{G}$, which was designed~\\citep{Zhang-etal2007} to characterize the linear growth of matter-density perturbations and which has been measured recently for moderate redshifts~\\citep{Reyes-etal2010}. In Section~\\ref{sec:Benchmark-LambdaCDM-results}, the same four observable quantities are calculated analytically for the standard $\\Lambda\\text{CDM}$--model universe~\\citep{Weinberg1972,Carroll-etal1992,SahniStarobinsky2000, PeeblesRatra2003,Perivolaropoulos2010} and may serve as benchmark results.  The numerical results for both homogeneous and inhomogeneous observable quantities of the QCD-scale modified-gravity universe are discussed in Section~\\ref{sec:Numerical-results} and compared with those of the $\\Lambda\\text{CDM}$ universe. Final comments are presented in Section~\\ref{sec:Discussion}.  Even though not necessary for a proper understanding of the present article, the reader may wish to have a look at Appendix~A of a recent review~\\citep{KV2011-review}, where the basic idea of $q$--theory is outlined.   %%\\newpage%%tmp \\setcounter{equation}{0} \\renewcommand{\\theequation}{\\arabic{section}.\\arabic{equation}} ",
        "conclusions": "\\label{sec:Discussion}  In this article, a simple empirical model has been proposed with a QCD-scale modified-gravity term \\eqref{eq:action-Sgrav0} and a single pressureless matter component (cold dark matter, CDM). This particular $f(R)$ modified-gravity model has been selected on physical grounds (Section~\\ref{subsec:Selection-criteria}), but is, in the end, solely used as an efficient way to describe the main aspects of the late evolution of the universe, having only two fundamental energy scales, $E_\\text{QCD}$ and $E_\\text{Planck}$, and a single dimensionless coupling constant, $\\eta$.   With the elementary-particle-physics value for the equilibrium gluon condensate $q_{0} \\equiv (E_\\text{QCD})^4 = (300\\;\\text{MeV})^4$, the measured age of the universe fixes the dimensionless coupling constant $\\eta$ of model \\eqref{eq:action-Sgrav0} to the value of approximately $2.4 \\times 10^{-4}$. As emphasized in \\citet{K2010-gluoncondcosm}, the effective coupling constant $\\eta$ may ultimately be calculated from QCD and general relativity; see also \\citet{Schuetzhold2002,Bjorken2004,UrbanZhitnitsky2010a, UrbanZhitnitsky2010b,Holdom2011} for further discussion of the possible relation of QCD and dark energy. The connection of the QCD-scale modified-gravity model with the $q$--theory approach to solving the main cosmological constant problem has already been mentioned in Section~\\ref{subsec:Selection-criteria}.  The $q_{0}$ and $\\eta$ values quoted in the previous paragraph result in a model postdiction for the Hubble constant \\eqref{eq:results-Hp}, which is within $10\\,\\%$ of the observed value. Theoretically, the last mathematical expression for the Hubble constant in the middle of \\eqref{eq:results-Hp} is quite remarkable, as it involves only Newton's constant $G_{N}$ and the cube of the energy scale $E_\\text{QCD}$. Independent of the $q_{0}$ and $\\eta$ values, there are further model predictions for the dimensionless quantities \\eqref{eq:3obs-QCDmodgrav} and \\eqref{eq:E_G-estimate}, which are again in the same ball park as the observed values. The modified-gravity model prediction of $-0.7$ for the quantity $\\overline{w}_{X}(\\tau_{p})$ as defined by \\eqref{eq:modgrav-wXbar-tp} [compared to the $\\Lambda\\text{CDM}$ value of $-1$] may perhaps provide a crucial test, as long as \\emph{independent} measurements of the present values of $H$, $d H/d\\tau$, and $\\rho_\\text{CDM}\\equiv \\rho_\\text{M2}$ can be obtained.  Moreover, model predictions for the redshift dependence of the gravity estimator $E_{G}$ have been given in Table~\\ref{tab-H-EG-results}. At the very largest length scales (comoving wavelengths of the order of hundred Mpc or more but still less than the horizon scale) and relatively low redshifts $z \\sim 0.75$, the QCD-modified-gravity results for $E_{G}$ differ by some $+10\\%$ from the  $\\Lambda$CDM--model values. The QCD-modified-gravity results for $E_{G}$ differ by some $-10\\%$ from the $\\Lambda$CDM--model values for $z \\gg 3$, but it is not yet clear how these redshifts can be probed observationally.  As it stands, the QCD-scale modified-gravity model \\eqref{eq:action-Sgrav0} gives a remarkable description of the main aspects of the late evolution of the universe. With many forthcoming ground-based experiments and space missions, future measurements of $\\overline{w}_{X}(\\tau_{p})$ and $E_{G}(z)$ may suggest alternative $f(R)$ functions, such as the one given by \\eqref{eq:f-ext}. Of course, if experiment indeed finds evidence for nonzero $f(R)$, it is up to theory to provide the proper understanding."
    },
    "1005/1005.0857_arXiv.txt": {
        "abstract": "Recent observational results for the masses and radii of some neutron stars are in contrast with typical observations and theoretical predictions for ``normal'' neutron stars. We propose that their unusual properties can be interpreted as the signature of a dark matter core inside them. This interpretation requires that the dark matter is made of some form of stable, long-living or in general non-annihilating particles, that can accumulate in the star \\cite{Sandin:2008db}. In the proposed scenario all mass-radius measurements can be explained with one nuclear matter equation of state and a dark core of varying relative size. This hypothesis will be challenged by forthcoming observations and could eventually be a useful tool for the determination of dark matter. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.2063_arXiv.txt": {
        "abstract": "We present a path-integral likelihood formalism that extends parameterized likelihood analyses to include continuous functions. The method finds the maximum likelihood point in function-space, and marginalizes over all possible functions, under the assumption of a Gaussian-distributed function-space. We apply our method to the problem of removing unknown systematic functions in two topical problems for dark energy research : scale-dependent galaxy bias in redshift surveys; and galaxy intrinsic alignments in cosmic shear surveys. We find that scale-dependent galaxy bias will degrade information on cosmological parameters unless the fractional variance in the bias function is known to $10\\%$. Measuring and removing intrinsic alignments from cosmic shear surveys with a flat-prior can reduce the dark energy Figure-of-Merit by $20\\%$, however provided that the scale and redshift-dependence is known to better than $10\\%$ with a Gaussian-prior, the dark energy Figure-of-Merit can be enhanced by a factor of two with no extra assumptions. ",
        "introduction": "The Standard Cosmological Model is a phenomenological model containing two unknown components, dark energy and dark matter, and two untested assumptions, general relativity on cosmic scales and cosmological inflation. To understand this model at a deeper level, by testing the evolution of dark energy, the nature of the dark matter, signatures of inflation and possible deviations from Einstein gravity, we need highly sensitive cosmological probes that must constrain an expanded set of cosmological parameters to (sub-)percent accuracy. In such a scenario the correct assessment of systematic effects is of critical importance.  To account for systematic effects it is common to use a parameterized description of the effect where each nuisance parameter is calibrated to some accuracy and removed. However, there are many cases where characterization of the systematic effect by a parameterized function is not well-justified, and in some cases the systematic is effectively an unknown function. The assumption of a given functional form can lead to a bias in measured cosmological parameters, or at least an underestimate of the contribution to errors due to uncertainty in the systematic function.  In Taylor \\& Kitching (2010) we described a new method which allowed the analysis of large-dimensional parameter-spaces composed of cosmological parameters and nuisance parameters describing residual systematics, in both calibrated (Gaussian prior) and self-calibrating (flat-prior) regimes. Here we present a new method that can account for systematics in a model-free way such that the impact of every possible function on the cosmological signal is accounted for. To do this we introduce a path-integral marginalization approach in which we maximize the likelihood and  marginalize over a space of functions rather than parameters. We construct a likelihood functional in parameter space that accounts for the impact of all known systematic functions. The function-space may be either weighted with a flat prior about a fiducial function, where the data itself is used to measure and remove the systematic effect, or if there is external data weighted with a Gaussian prior, with a variance that is zero when the systematic function is known.  We derive an expression for the Fisher matrix that accounts for functional marginalization, and apply this method to two test cases in cosmology in which the cosmological signal is contaminated by a systematic effect with a poorly understood functional behaviour. The first is scale-dependent galaxy bias (e.g.,  Peacock \\& Smith, 2000; Conway et al., 2005; Hamann et al., 2008; Cresswell \\& Percival, 2009) in redshift surveys, an unknown function that relates the underlying dark matter distribution to the galaxy distribution. The second is cosmic shear intrinsic alignments (e.g. Brown et al., 2002 ; Crittenden et al , 2001, Catelan et al,, 2001; Heymans et al., 2004; Hirata \\& Seljak, 2004; Kitching et al., 2007; Bridle \\& King, 2007; Joachimi \\& Bridle, 2009; Joachimi \\& Schneider, 2008, 2009; Mandelbaum et al., 2009; Schneider \\& Bridle, 2009; Kirk et al., 2010) in weak lensing surveys, a contaminant with a complex astrophysical function that relates the local density to the ellipticity of background galaxy images. In both cases we show how cosmological parameter constraints are expected to degrade when we marginalize over all systematic functions.  This work is an extension of the methods of  Taylor \\& Kitching (2010) in which a formalism for analytic marginalization was developed for a finite number of parameters, and is related to the Functional Form-Filling method of Kitching et al. (2009) in which a numerical approach to assessing the effects of unknown functions was presented.  This paper is organized as follows, in Section \\ref{Path Integral Marginalization} we introduce path integral marginalization, in Sections \\ref{Scale-dependent, stochastic galaxy bias} and \\ref{Cosmic Shear Intrinsic Alignments} we apply the method to galaxy bias and cosmic shear intrinsic alignments respectively. In Section \\ref{Conclusions} we present conclusions. Throughout we have put technical detail in a series of Appendices. ",
        "conclusions": "\\label{Conclusions}  Path-integral marginalization can account for the impact of measuring and marginalizing over unknown functions on cosmological parameter estimation. We have constructed a likelihood function in cosmological parameter space that takes into account marginalization over unknown functions, assuming either a flat-prior, were we use information in the data itself to measure and remove the effect, or a Gaussian-prior, where we use external data to constrain the nuisance functions. We assume the distribution in function-space is well approximated by a Gaussian functional, but note that the non-Gaussian case can be treated with numerical methods.  We use the likelihood to derive an expression for the cosmological parameter Fisher information matrix in the case that all functions are marginalized over. We have applied this to the case of a scale-dependent galaxy bias function, relating the galaxy power spectrum to the underlying matter power spectrum. As a scale-dependent bias is completely degenerate with the matter power spectrum we assume a Gaussian-prior with variable scatter. We show that the marginal errors of $\\Omega_m$, $w_0$ and $w_a$ rapidly increase as the uncertainty in the fractional power of a scale-dependent bias increases above $P_b(k)/b^2(k)\\gs 10^{-1}$, implying we need to constrain any scale-dependent bias to a few percent accuracy. We did not explore redshift-dependent bias models here, but the marginalized Fisher matrix implies a similar accuracy is needed. In addition we did not explore stochastic bias models, which will add a further requirement. There are clear ways to constrain scale-dependent bias from redshift-space distortions, which also probe the velocity field and so are bias-free; by the galaxy bispectrum; or by combining galaxy redshift surveys with weak lensing probes of the dark matter (e.g. van Waerbeke, 1998; Bernstein, 2009; Matarrese et al., 1997; Verde et al., 1998; Guo \\& Jing, 2009).  Our second example is in the removal of Intrinsic Alignments effects, in particular the non-local GI effect, on cosmic shear measurements. In addition to nulling (Joachimi \\& Schneider, 2008, 2009) and modelling (Bridle \\& King, 2007; Kitching et al., 2008; Joachimi \\& Bridle, 2009), path-integral marginalization represents a third way to remove and assess the impact of this effect. Path-integral marginalization has the advantage that it is model-free, and introduces no biases in cosmological parameter estimates. For a flat-prior on the ellipticity-density relation, the dark energy FoM drops from a fiducial ${\\rm FoM}=130$, to ${\\rm FoM}\\approx 100$, when we allow the GI term to be completely free in both angular and radial dependence and measured by the data.  If we assume the GI ellipticity-density relation is accurately constrained from external data (e.g. Mandelbaum et al., 2009) or complimentary information (e.g. shear-position information; Bernstein, 2009) with a Gaussian-prior, and allow the weak lensing effect in the GI term to depend on cosmology, we find that the dark energy constraint is enhanced from FoM$=130$ to FoM$=300$. The FoM is degraded as the accuracy drops below $\\sigma(P_{\\gamma \\delta})/P_{\\gamma\\delta}\\approx 10^{-3}$. The assumption of a non-linear alignment model (e.g Hirata \\& Seljak, 2004; Bridle \\& King, 2007) relaxes this requirement to $\\approx 10^{-1}$.  Our results imply the GI effect in weak lensing can be removed by path-integral marginalization and that cosmic shear can still constrain dark energy to high precision. In addition, with no extra assumptions, we have shown that the GI effect can itself be used to probe cosmology through its dependence on weak lensing. By constraining the GI ellipticity-density dependence, using external data, we can use the weak lensing dependence in GI to increase our knowledge of dark energy.  The path-integral marginalization approach presented here is a powerful and general method to assess, control and remove the effects of systematics which cannot be simply modelled as parametric functions.   \\vspace{0.5cm}  \\noindent{\\em Acknowledgements:} We thank Fergus Simpson, Adam Amara, Benjamin Joachimi, Andrew Liddle, John Peacock and Alan Heavens for useful discussions. We also thank the DUEL network (MRTN-CT-2006-036133) for supporting this work. TDK is supported by STFC rolling grant number RA0888.  \\vspace{-0.5cm}"
    },
    "1005/1005.0250_arXiv.txt": {
        "abstract": "Ultrahigh energy neutrons and pions are likely to be produced in particle interactions inside cosmic ray sources and subsequently decay to neutrinos and other secondary particles [$\\pi^{\\pm}\\rightarrow \\mu^{\\pm}\\nu_{\\mu}(\\bar \\nu_{\\mu}),\\mu^{\\pm}\\rightarrow e^{\\pm}\\bar\\nu_{\\mu}(\\nu_{\\mu})\\nu_e(\\bar\\nu_e)$]. In high magnetic fields of the cosmic acceleration sites, the ultrahigh-energy-charged particles may lose energy significantly due to synchrotron radiation before decay. We show that for gamma ray bursts in the internal shock model the flux of very high-energy antineutrinos ($\\bar \\nu_e$) produced from decaying ultrahigh-energy neutrons can be more than the total neutrino flux produced in pion decay depending on the values of their Lorentz factors, luminosities and variability times. ",
        "introduction": "We discuss the photopion and neutron decay channels of high-energy neutrino production inside GRBs. The high-energy neutrino flux generated in $p\\gamma$ interactions through pion decay from individual GRBs in the internal shock model has been calculated earlier \\cite{nay1,guetta}. In \\cite{nay2} the authors have derived the neutrino flux from individual GRBs using the isotropic photon energy typically observed by {\\it Swift} in the energy range of 1 KeV to 10 MeV and folded the neutrino flux with the distribution functions to incorporate the fluctuations in the values of the GRB parameters. The diffuse neutrino flux from all high and low- luminosity GRBs have been calculated in \\cite{nay2} considering distribution functions for the GRB parameters. GRBs were assumed to follow star formation rate and the rate of low luminosity GRBs was taken to be 500 times higher than high luminosity GRBs. In this paper we have calculated the neutrino flux from individual GRBs by expressing it in terms of the observed isotropic photon energy as discussed in \\cite{nay2} assuming certain values for the luminosity, Lorentz factor and variability time of a GRB. We have not considered the distribution in the values of the GRB parameters which is important in case of the diffuse flux from all GRBs. The photon energy spectrum from a GRB can be expressed as a broken-power law with break at $\\epsilon_{br}$ in the source rest frame. \\beq \\frac{dn_{\\gamma}}{d\\epsilon_{\\gamma}} =A \\left\\{ \\begin{array}{l@{\\quad \\quad}l} {\\epsilon_{\\gamma}}^{-\\gamma_1} & \\epsilon_{\\gamma}<\\epsilon_{br}\\\\{\\epsilon_{br}}^{\\gamma_2-\\gamma_1} {\\epsilon_{\\gamma}}^{-\\gamma_2} & \\epsilon_{\\gamma}>\\epsilon_{br} \\end{array}\\right. \\eeq  These photons are interacting with shock accelerated protons to produce charged and neutral pions. The charged pions subsequently decay to muons and neutrinos. Muons also decay to neutrinos. The minimum energy of the protons interacting with photons of energy $\\epsilon_{br}$ is \\beq E_{pb}=1.3\\times10^{7}\\Gamma_{300}^2(\\epsilon_{br,MeV})^{-1} {\\rm GeV}. \\eeq In $p\\gamma$ interactions both $\\pi^{0}$ and $\\pi^{+}$ can be produced with equal probabilities. $\\pi^{+}$ gets, on average $20\\%$ of the proton energy and if the final-state leptons share the pion energy equally then each neutrino carries $5\\%$ of the initial proton energy. The first break in energy in the neutrino spectrum, $E_{\\nu b}$ is due to the break in the photon spectrum at $\\epsilon_{br}$. \\beq E_{\\nu b}=6.5\\times10^5\\frac{\\Gamma^2_{300}} {\\epsilon_{br,MeV}} {\\rm GeV} \\label{Eb1} \\eeq The fire ball Lorentz factor $\\Gamma_{300}= \\Gamma/{300}$, photon luminosity $L_{\\gamma,51}=L_{\\gamma}/(10^{51} ergs \\, /sec)$, variability time $t_{v,-3}=t_v/(10^{-3})sec$ are the important parameters in our calculations. The total energy to be emitted by neutrinos of energy $E_{\\nu}$ can be expressed as, \\beq E_{\\nu}^2\\frac{dN_{\\nu}(E_{\\nu})}{dE_{\\nu}}\\approx \\frac{3f_{\\pi}}{8}\\frac{1}{\\kappa}\\frac{(1-\\epsilon_e-\\epsilon_B)} {\\epsilon_e}E_{\\gamma}^{iso} \\label{spect} \\eeq where $E_{\\gamma}^{iso}$ is the total isotropic energy of the emitted gamma-ray photons in the energy range of 1 KeV to 10 MeV; it is available from observations. $\\epsilon_e,\\epsilon_B\\sim 0.3$ are the energy fractions carried by electrons and the magnetic field respectively. $f_{\\pi}$ is the fraction of the energy of a proton going to pion production inside a GRB fireball. The relativistic electrons produce the photons by synchrotron radiation, so four orders of magnitude in photon energy corresponds to two orders of magnitude in the energy of the radiating charged leptons. We have assumed that the electrons are radiating all their kinetic energy as observed photon energy $E_{\\gamma}^{iso}$. The normalization constant of the power law electron flux is proportional to $E_{\\gamma}^{iso}/\\kappa$. The differential spectral index of the electron spectrum is assumed to be $-2.5$ which corresponds to $\\kappa=1.8$. The total energy carried by the differential flux of neutrinos has been approximately related to the total energy carried by the differential flux of electrons in Eq.(\\ref{spect}). The first break [Eq.(\\ref{Eb1})] in the neutrino spectrum due to the break energy in the photon spectrum is contained in the expression of $f_{\\pi}$ \\cite{nay2}. The charged pions and muons lose energy due to synchrotron radiations in high magnetic fields inside the sources before decaying to neutrinos. As a result, the neutrino flux is depleted at very high energy. The neutrino flux from photopion decay in the source rest frame is \\beq E_{\\nu}^2\\frac{dN_{\\nu}(E_{\\nu})}{dE_{\\nu}}\\approx\\frac{3f_{\\pi}} {14.4}\\frac{(1-\\epsilon_e-\\epsilon_B)}{\\epsilon_e}E_{\\gamma}^{iso} \\left\\{\\begin{array}{l@{\\quad \\quad}l} 1 & E_{\\nu}<E_{\\nu}^{s}\\\\ (\\frac{E_{\\nu}}{E_{\\nu}^s})^{-2}& E_{\\nu}>E_{\\nu}^{s} \\end{array} \\right. \\eeq Pion-cooling energy is 10 times higher than muon-cooling energy. We would be overestimating the total neutrino flux if we use the pion-cooling energy to derive the second break energy in the neutrino spectrum. The muon-cooling break $E_{\\nu}^s$ can be expressed as a function of GRB parameters. \\beq E_{\\nu}^{s}=2.56\\times10^{6}\\epsilon_e^{1/2}\\epsilon_B^{-1/2} L_{\\gamma,51}^{-1/2}\\Gamma_{300}^4t_{v,-3}{\\rm GeV} \\eeq  The fraction of fireball proton energy lost to neutron production is, \\beq f_n (E_p)= 0.8f_0 \\left\\{\\begin{array}{l@{\\quad\\quad}l} \\frac{1.34^{\\gamma_1-1}}{\\gamma_1+1}({E_p}/{E_p}^b)^{\\gamma_1-1}& {E_p>E_p^b}\\\\ \\frac{1.34^{\\gamma_2-1}}{\\gamma_2+1}(E_p/{E_p}^b)^{\\gamma_2-1}& {E_{p}< E_{p}^b} \\end{array} \\right. \\label{fn} \\eeq where $f_0=\\frac{0.45L_{\\gamma,51}}{\\Gamma_{300}^4 \\, t_{v,-3}{\\epsilon_{br,MeV}}}\\frac{1}{\\big [\\frac{1}{\\gamma_2-2}-\\frac{1}{\\gamma_1-2}\\big ]}$, details are discussed in \\cite{nay2}. The average value of the fraction of energy transferred from the proton to the neutron is assumed to be $\\left< x_{p\\rightarrow n} \\right>= 0.8$. The neutron spectrum can be expressed in terms of $E_{\\gamma}^{iso}$ and the equipartition parameters $\\epsilon_e$, $\\epsilon_B$. \\beq E_n^2\\frac{dN_n}{dE_n}=\\frac{f_n}{2\\kappa}\\frac{1-\\epsilon_e-\\epsilon_B}{\\epsilon_e}{E_{\\gamma}^{iso}} \\eeq $\\frac{1}{2}$ corresponds to the assumption that charged and neutral pions are produced with equal probabilities. The GRB is assumed to be at a redshift of 0.8 in all our calculations. Duration of the burst is assumed to be 5 sec. The break in the photon spectrum is assumed to be at $\\epsilon_{br}=0.3MeV$ in the source rest frame. We have assumed the spectral indices of the photon flux $\\gamma_1$=1.5 and $\\gamma_2=2.2$. The ultrahigh energy neutrons decay ($n\\,\\rightarrow p +e^-+\\overline{\\nu}_e$) to $\\bar\\nu_e$. Their rest frame lifetime is $\\overline{\\tau}_n\\,=\\,886$ seconds and decay mean free path $10(E_n/EeV)$ kpc. The maximum energy of the neutrons producing $\\bar\\nu_e$ depends on the distance $D_s$ of the source. If the source is at a large distance then all the neutrons are expected to decay during propagation. The $\\bar\\nu_e$ flux expected on earth at energy $E_{\\bar{\\nu}}$ from decaying neutrons ( mass $m_n$) of flux $({dN_n}/{dE_n})$ \\cite{anc2} is \\beq \\frac{dN_{\\bar{\\nu}}}{dE_{\\bar{\\nu}}}(E_{\\bar{\\nu}})=\\frac{m_n}{2\\,\\epsilon_0}\\int_{\\frac{m_n\\,E_{\\bar{\\nu}}}{2\\,\\epsilon_0}}^{E_{n,max}}\\frac{dE_n}{E_n}\\,\\frac{dN_n}{dE_n}\\left(1-e^{-\\frac{D_s\\,m_n}{E_n\\,\\overline{\\tau}_n}}\\right)\\, \\eeq where $\\epsilon_0$ is the mean energy of $\\bar\\nu_e$ in the neutron rest frame. The maximum energy of the shock accelerated protons $E_{p,max}$ in the GRB fireball can be $10^{12}$ GeV. The maximum energy of the secondary neutrons is $E_{n,max}=0.8E_{p,max}$, assuming $80\\%$ of the proton's energy goes to a neutron in $p\\gamma$ interactions. The maximum energy of the antineutrinos is $Q/m_n=0.001 $ times  the maximum energy of the decaying neutrons, where $Q=m_n-m_p-m_e$. The observed antineutrino/neutrino flux on earth is \\beq \\frac{dN_{\\nu}^{ob}(E_{\\nu}^{ob})}{dE_{\\nu}^{ob}}=\\frac{dN_{\\nu}(E_{\\nu})}{dE_{\\nu}}\\frac{1+z}{4\\pi D_s^2} \\eeq Figures 1-3 show our calculated neutrino fluxes from photopion decay (dashed line) and neutron decay (solid line) in the observer's frame. In Fig.1. the Lorentz factor has been varied from 100 to 500 assuming luminosity $10^{51}$erg/sec and variability time 10ms. For small values of Lorentz factor  and variability time the neutron decay channel can become more important than the photo-pion decay channel at very high energy, as shown in Figs.1 and 3. respectively. Figure 2. shows for high values of luminosity this effect is more significant. The number of neutrino events expected from a GRB in IceCube is very low. The neutron decay channel can give only $10^{-5}$ muon events in $Km^2$ area above 1 TeV energy from a GRB at redshift 0.8, with luminosity $L_{\\gamma}=10^{53}$erg/sec, Lorentz factor $\\Gamma=300$, duration 5sec, and variablity time $t_v=10$ms. The GRB has been assumed to occur at a zenith angle of $120^{o}$ with respect to an observer at the south pole. The number of total neutrino events contained inside $2km^3$ volume of ice is expected to be $10^{-5}$ from the neutron decay channel and $10^{-4}$ from the pion decay channel. Above $10^{16}$ or $10^{17}$ eV we expect the neutron decay channel to dominate over the pion decay channel depending on the values of the GRB parameters. Neutrino telescopes like ANITA  may be useful for detecting the neutrino fluxes in this energy range. ",
        "conclusions": "Multimessenger approach to explore the cosmic-ray universe has gained much importance due to the successful operations of the cosmic-ray, gamma-ray and neutrino detectors. In this paper we have discussed that the antineutrino flux produced in decay of ultrahigh-energy neutrons may become more important than the neutrino flux from photopion decay depending on the luminosity , Lorentz boost factor, varibility time of a GRB at very high energy. In future one may consider other cosmic acceleration sites with high magnetic fields to explore the importance of the neutron decay channel at very high energy."
    },
    "1005/1005.1640_arXiv.txt": {
        "abstract": "The Parker model for coronal heating is investigated through a  high resolution simulation. An inertial range is resolved where fluctuating magnetic energy $E_{_M}(k_{\\perp}) \\propto k_{\\perp}^{-2.7}$ exceeds kinetic energy $E_{_K}(k_{\\perp}) \\propto k_{\\perp}^{-0.6}$. Increments scale as $\\delta b_{_{\\!\\ell}} \\simeq  \\ell^{-0.85}$ and $\\delta u_{_{\\!\\ell}} \\simeq  \\ell^{+0.2}$ with velocity increasing at small scales, indicating that magnetic reconnection plays a prime role in this turbulent system. We show that spectral energy transport is akin to standard magnetohydrodynamic (MHD) turbulence even for a system of reconnecting current sheets sustained by the boundary. In this new MHD turbulent cascade, kinetic energy flows are negligible while cross-field flows are enhanced, and through a series of ``reflections'' between the two fields, cascade more than half of the total spectral energy flow. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.0470_arXiv.txt": {
        "abstract": "It is believed that the local massive black holes were dominantly grown up through accretion during quasar phases, while a fraction of the local black hole mass was accumulated through accreting gases at very low rates. We derive the black hole mass density as a function of redshift with the bolometric luminosity function of active galactic nuclei (AGN) assuming that massive black holes grew via accreting the circumnuclear gases, in which the derived black hole mass density is required to match the measured local black hole mass density at $z=0$. Advection dominated accretion flows (ADAFs) are supposed to present in low luminosity active galactic nuclei (AGNs)/normal galaxies, which are very hot and radiate mostly in the hard X-ray band. Most of the X-ray background (XRB) is contributed by bright AGNs, and a variety of AGN population synthesis models were developed to model the observed XRB in the last two decades. Based on our derived black hole mass density, we calculate the contribution to the XRB from the ADAFs in faint AGNs/normal galaxies with a given Eddington ratio distribution, which is mostly in hard X-ray energy band with an energy peak at $\\sim 200$~keV. The growth of massive black holes during ADAF phase can therefore be constrained with the observed XRB. Combining an AGN population synthesis model with our results, {we find that the fitting on the observed XRB, especially at hard X-ray energy band with $\\ga 100$~keV, is improved provided the contribution of the ADAFs in low-luminosity AGNs/normal galaxies is properly included.} It is found that less than $\\sim$15 per cent of local massive black hole mass density was accreted during ADAF phases. {We suggest that more accurate measurements of the XRB in the energy band with $\\ga 100$~keV in the future may help constrain the growth of massive black holes at their late stage.} We also calculate their contribution to the extragalactic $\\gamma$-ray background (EGRB), and find that less than $\\sim1\\%$ of the observed EGRB is contributed by the ADAFs in these faint sources. ",
        "introduction": "It is well believed that almost all galaxies contain massive black holes at their centers, and a tight correlation was revealed between central massive black hole mass and the velocity dispersion of the galaxy \\citep[][]{f00,g00}. The growth of massive black holes at the centers of galaxies may probably be linked to accretion processes \\citep{s82}. \\citet{y02} estimated the black hole masses from the stellar velocity dispersions of galaxies measured by the Sloan Digital Sky Survey (SDSS) using the empirical relation between black hole mass and the velocity dispersion, and the local black hole mass density was derived. They further calculated the black hole mass density accreted during optical bright quasar phases using an optical quasar luminosity function (LF), and found that the accreted mass density is consistent with the local black hole mass density estimated from the velocity dispersions, if a radiative efficiency $\\sim 0.1$ is adopted for quasars \\citep*[also see][]{ma04,shankar04,h07,shankar09}. This implies that the growth of massive black holes through accretion during optically bright quasar phases may probably be important, if a radiative efficiency $\\sim 0.1$ is adopted. If the massive black holes are spinning rapidly, their radiative efficiency can be higher than that they adopted, the black hole mass density accreted during quasar phases would be lower than the measured local black hole mass density \\citep*[e.g.,][]{c07,cl08,l10}. Thus, one cannot neglect the contribution from accretion at low rates to the growth of massive black holes, provided the duration of accretion at low rates is as long as Hubble timescale.  The cosmological X-ray background (XRB) is mostly contributed by AGNs \\citep{h98,s98}, which can be used to constrain massive black hole accretion history \\citep*[e.g.,][]{e02}. In the most popular synthesis models of the XRB based on the unification schemes for AGNs, the cosmological XRB contributed by Compton-thin AGNs can account for $\\sim 80$\\% of the observed XRB \\citep*[e.g.,][]{u03,g07}. The residual XRB can be explained provided the same number of Compton-thick AGNs with $\\log N_{\\rm H}=24-25$ as those with $\\log N_{\\rm H}=23-24$ is included \\citep{u03}. \\citet{d97} have alternatively proposed that the hard XRB above 10 keV may be dominated by the thermal bremsstrahlung emission from the advection dominated accretion flows (ADAFs) in low-luminosity AGNs. {Due to the difficulties on detecting Compton-thick AGNs, the space number density of Compton-thick AGNs is still quite debated \\citep*[e.g.,][]{u03,tu05,g07,tuv09}. Based on the surveys with \\textit{Swift} and \\textit{INTEGRAL} satellites, \\citet{tuv09} suggested that the space number density of Compton-thick AGNs may be significantly lower than those adopted in the previous AGN population synthesis models to explain the observed XRB.} Due to the uncertainty of the number density of Compton-thick AGNs, the residual XRB could be attributed to both the Compton-thick AGNs and the ADAFs in faint AGNs and normal galaxies. Therefore, the observed XRB can be used to constrain massive black hole accretion history at their late stage when they are accreting at low rates \\citep{c05,c07}. \\citet{c07} calculated the contribution to the XRB from the ADAFs in faint AGNs/normal galaxies, and compared it with the residual XRB. Their results showed that less than $\\sim$5 per cent of local massive black hole mass density was accreted during ADAF phases, otherwise the XRB contributed from ADAFs will surpass the observed residual XRB even if no Compton-thick AGNs are included.  For simplicity, they adopted an average mass accretion rate for faint AGNs/normal galaxies in calculating the contribution of ADAFs to the XRB. The radiative efficiency of standard thin accretion disks in bright AGNs does not vary with mass accretion rate. However, the faint AGNs/normal galaxies may probably contain ADAFs, of which the radiative efficiencies vary with mass accretion rate \\citep{n95}. The average radiative efficiency for a population of sources containing ADAFs accreting at different rates can be calculated by weighing over the distribution of mass accretion rate. The Eddington ratios of AGNs spread over several orders of magnitude \\citep*[e.g.,][]{h02,h06,cx07}. The Eddington ratio distribution for accreting massive black hole holes derived from observations exhibits nearly a power-law distribution with an exponential cutoff at high Eddington ratio \\citep{m08,h09,kh09}. In this paper, we will adopt the Eddington ratio distribution of AGNs derived by \\citet{h09} to calculate the contribution of faint AGNs/normal galaxies to the XRB. Compared with the residual XRB, the constraints on the fraction of local black hole mass accreted in ADAF phases are derived. The cosmological parameters $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0=70~ {\\rm km~s^{-1}~Mpc^{-1}}$ have been adopted in this work. ",
        "conclusions": "Unlike standard thin accretion disks, the radiative efficiency of ADAFs increases with mass accretion rate $\\dot{m}$ \\citep{n95}. Thus, the average radiative efficiency for a population of ADAFs is available only if the distribution of mass accretion rate $\\dot{m}$ is known. The Eddington ratio distribution for AGNs was extensively explored by different authors \\citep*[e.g.,][]{m08,h09,kh09}. These works suggested that more sources are accreting at lower rates, and the Eddington ratios for AGNs can be well described by a power law distribution \\citep{h09}.  This implies that the average radiative efficiency for the ADAFs with this Eddington ratio distribution can be lower than that for an ADAF accreting at the average mass rate, because the radiative efficiency of ADAFs increases with mass accretion rate $\\dot{m}$ \\citep{n95}. \\citet{c07} calculated the contribution of ADAFs to the XRB, and the constraint on the fraction of black hole mass accreted in ADAF phases was derived by comparing with the residual XRB. An average mass accretion rate is adopted in \\citet{c07}'s calculations for simplicity, and therefore the upper limit on the mass growth of black holes in ADAF phases is under-estimated. In this work, we improve the calculations by \\citet{c07} by including the Eddington ratio distribution for AGNs in our calculations.  Using the method described in \\S 2, the active, inactive and total black hole mass densities as functions of redshift are calculated by tuning the average radiative efficiency $\\epsilon$ for active galaxies to let the total black mass density match the measured local black hole mass density at $z=0$ (see Fig. \\ref{fig_zrho}). It is found that most of the local black hole mass was accumulated during $z\\la 2$, and fraction of the local black hole mass accreted during ADAF phases is determined by the average mass accretion rate $\\dot{m}_{\\rm inact}^{\\rm aver}$, i.e., the distribution of accretion rates for inactive galaxies. We find that the average radiative efficiency $\\simeq 0.134$ is required for active galaxies, which corresponds to $a\\approx 0.9$. This is consistent with that derived in \\citet{c07} and \\citet{e02}.   {In the calculation of the black hole mass densities, we need to know the value of the average Eddington ratio for the active galaxies $\\lambda_{\\rm act}^{\\rm aver}$, which is still uncertain. \\citet{md04} estimated that the average accretion rate $\\dot{m}^{\\rm aver}_{\\rm act}$ varies from 0.1 at $z\\sim0.2$ to 0.4 at $z\\sim2$ from a large sample of SDSS quasars. \\citet{k06} estimated \\texttt{}the black hole masses and Eddington ratios for a sample of luminous AGNs with $0.3<z<4$, and found that their average Eddington ratio is $\\simeq 0.25$. We adopt $\\lambda_{\\rm act}^{\\rm aver}=0.25$ in this work in the calculation of active black hole mass density $\\rho_{\\rm bh}^{\\rm act}(z)$ and then the inactive black hole mass density $\\rho_{\\rm bh}^{\\rm inact}(z)$. The derived active black hole mass density $\\rho_{\\rm bh}^{\\rm act}(z)$ is proportional to $1/\\lambda_{\\rm act}^{\\rm aver}$ (see Eq. \\ref{bhmdensact}). We find that the derived $\\rho_{\\rm bh}^{\\rm act}(z)\\ll \\rho_{\\rm bh}^{\\rm inact}(z)$ for $z\\la 3$ (see Fig. \\ref{fig_zrho}). The inactive black hole mass density $\\rho_{\\rm bh}^{\\rm inact}(z)$ is calculated by subtracting the active black hole mass density $\\rho_{\\rm bh}^{\\rm act}(z)$ from the total black hole mass density $\\rho_{\\rm bh}(z)$ at redshift $z$ (see Eq. \\ref{rhobhinact}), and $\\rho_{\\rm bh}(z)$ is derived with QLF as required to match the measured local black hole mass density $\\rho_{\\rm bh}^{\\rm local}$ at $z=0$, which implies that the derived $\\rho_{\\rm bh}^{\\rm inact}(z)$ is almost insensitive to the value of $\\lambda_{\\rm act}^{\\rm aver}$ for $z\\la3$. The XRB and growth of massive black holes are mainly contributed by the accretion in the sources at low redshifts, which implies that our main conclusions on the growth of massive black holes at their late stage will not be affected if the value of $\\lambda_{\\rm act}^{\\rm aver}$ does not deviate much from the value adopted in this work. For the average radiative efficiency $\\epsilon=0.134$ derived in this work, the average mass accretion rate for active galaxies $\\dot{m}_{\\rm act}^{\\rm aver}=0.187$ corresponding to $\\lambda_{\\rm act}^{\\rm aver}=0.25$. The black hole mass densities can be calculated alternatively by using the Eddington ratio distribution given by Eq. \\ref{mdot_dis}. This distribution shows that most active galaxies have mass accretion rates $\\dot{m}$ close to the critical one $\\dot{m}_{\\rm crit}=0.01$ with an average $\\dot{m}_{\\rm act}^{\\rm aver}\\sim 0.1$ for $\\kappa_{\\rm m}=0.3$. As discussed above, our main conclusions will not be altered either with this value or a distribution instead of a single average one. The contribution of all bright AGNs to the XRB can be directly calculated with the QLF, which is independent of the Eddington ratio distribution. \\citet{h07}'s calculation showed that the observed XRB can be roughly reproduced with their QLF. }    Our calculations of the XRB contributed by ADAFs in inactive galaxies are sensitive to the spectra of ADAFs, which are described by several parameters, i.e., magnetic field strength relative to gas pressure $\\beta$, the fraction of energy directly heating the electrons $\\delta$, and the viscosity parameter $\\alpha$. In \\citet{c07}'s work, the theoretical/observational constraints on the values of these parameters were discussed, which leads to narrow ranges for the values of these parameters. The X-ray/$\\gamma$-ray spectra of ADAFs are almost independent of the value of $\\beta$ (magnetic field strength), which only affects the spectra in radio bands. As we are focusing on the hard X-ray and $\\gamma$-ray energy bands, we find that the spectra of ADAFs depend most sensitively on the value of $\\delta$, because the electron temperature in the ADAF is sensitively affected by the energy directly heating the electrons. Besides a conventional value of $\\delta=0.1$ adopted in our calculations, we also carry out the calculations with $\\delta=0.01$ for comparison. In this case, the radiative efficiency of ADAFs is lower than that for $\\delta=0.1$, as the heating of electrons is significantly suppressed (see Figs. \\ref{fig_nulnu_delta0p1}-\\ref{fig_lambda_mdot}).  Figure \\ref{fig_lambda_mdot} shows how the Eddington ratios $L_{\\rm bol}/L_{\\rm Edd}$ vary with mass accretion rate $\\dot{m}$ for different values of $\\delta$ and black hole { spin parameter} $a$. We find that the Eddington ratio varies with $\\dot{m}$ roughly as $L_{\\rm bol}/L_{\\rm Edd}\\propto \\dot{m}^s$, where $s\\simeq 1.2-2.1$ depending on the values of parameters adopted. {For comparison, we plot several power-law lines $L_{\\rm bol}/L_{\\rm Edd}\\propto \\dot{m}^s$, $s=1$, i.e., constant radiative efficiency $\\epsilon=0.1$ (black dash-dotted line), $s=1.2$ and $2.1$ in the same figure.} The result $s\\simeq 1.2-2.1$ means that the radiative efficiencies of ADAFs increases with $\\dot{m}$. Our results are consistent with those adopted in the previous works \\citep{n95,j05,h07}. It is not surprising that the radiative efficiency increases with $\\delta$, and the bolometric luminosity $L_{\\rm bol}$ is higher for rapidly spinning black holes provided the same accretion rate $\\dot{m}$ is adopted.  Combining the derived inactive black hole mass density $\\rho_{\\rm bh}^{\\rm inact}(z)$ and the average spectra of a population of ADAFs with the mass accretion rate distribution (Eq. \\ref{mdot_dis}), the contribution of the ADAFs in all inactive galaxies to the XRB can be calculated. The results are plotted in Figs. \\ref{fig_ele_x_delta0p1} and \\ref{fig_ele_x_delta0p01}, which are compared with the observed XRB. The contribution of bright AGNs to the XRB was estimated in many previous works \\citep*[e.g.,][]{u03,tu05,g07,tuv09}. We compare the sum of the contribution of bright AGNs to XRB estimated by \\citet{tuv09} and that of the ADAFs in all inactive galaxies calculated in this work with the observed XRB. It is found that $\\kappa_{\\rm m}\\ga 0.3$ for $\\delta=0.1$ (or $\\kappa_{\\rm m}\\ga 0.2$ for $\\delta=0.01$) is required in order to let the calculated XRB not surpass the observed XRB (see Figs. \\ref{fig_ele_x_delta0p1} and \\ref{fig_ele_x_delta0p01}). {The dependence of required value of $\\kappa_{\\rm m}$ on parameter $\\delta$ can be understood that, a smaller $\\delta$ corresponds to a lower radiative efficiency, while a smaller $\\kappa_{\\rm m}$ represents a flatter dimensionless mass accretion rate distribution which has relatively more high- $\\dot{m}(\\approx10^{-3}-10^{-2})$ sources. The average mass accretion rate for inactive black holes is relatively high for a distribution with a small $\\kappa_{\\rm m}$, which also corresponds to a relatively high average radiative efficiency, because the radiative efficiency of an ADAF increases with $\\dot{m}$ (see Fig. \\ref{fig_lambda_mdot}). This implies that the distribution with a too flat mass accretion distribution (i.e., small $\\kappa_{\\rm m}$) will over-produce sources accreting at rates close to $\\dot{m}_{\\rm max}$, which will radiates too much to the XRB. On the other hand, the value of $\\kappa_{m}$ can also be constrained by $\\kappa$ and $s$ according to the relation $L_{\\rm bol}/L_{\\rm Edd}\\propto \\dot{m}^s$ and equations (\\ref{edd_ratio_dis}) and (\\ref{mdot_dis}). \\citet{h09} suggested $\\kappa\\approx 0.3-0.8$, and our calculation predicts $s\\simeq 1.2-2.1$ for different parameters, then we can expect $\\kappa_{\\rm m}\\approx \\kappa s\\approx 0.36-1.67$, which is roughly consistent with the constraints from comparison with the XRB in this work. Our results implies that the XRB contributed by inactive black holes is dominated by the radiation from high-$\\dot{m}$ sources, rather than low-$\\dot{m}$ sources. }  {The energy peak of the contribution of the ADAFs in inactive galaxies to the XRB is around $\\sim 100-200$~keV depending on the values of $\\delta$, which accounts for $\\sim15-20$\\% of the XRB at these energy peaks (see Figs. \\ref{fig_ele_x_delta0p1} and \\ref{fig_ele_x_delta0p01}). The \\textit{Swift} measurements on the XRB have estimated errors of $\\sim 3$\\% \\citep*{a08,tuv09}. Thus, we suggest that more accurate measurements of the XRB in the energy band with $\\ga 100$~keV in the future may constrain the growth of massive black holes at their late stage more precisely. We note that the peak energy of the contribution of Compton-thick AGNs is $\\sim 30$~keV, while it is $\\sim100-200$~keV for that of ADAFs in inactive massive black holes. This implies that the constraints on the growth of massive black holes at their late stage from the XRB have hardly been affected by the possible uncertainty of the space number density of Compton-thick AGNs. }  The values of $\\kappa_{\\rm m}$=0.2 and 0.3 correspond to average mass accretion rates of inactive galaxies $\\dot{m}_{\\rm inact}^{\\rm aver}=8.35\\times10^{-4}$ and $6.12\\times10^{-4}$, respectively (see Fig. \\ref{fig_kdotminact}). This means that the mass fraction of local black holes grown in ADAF phases should be less than $10.9\\%$ (for $\\delta=0.1$) and $14.5\\%$ (for $\\delta=0.01$) (see Fig. \\ref{fig_zrho}). Our result is about twice of that ($\\sim 5\\%$) for $\\delta=0.1$ given in \\citet{c07}. The discrepancy is mainly attributed to two factors. The first one is that a single average mass accretion rate adopted in \\citet{c07} for the calculations of ADAF spectra and their contribution to the XRB, which over-estimated the average radiative efficiency for ADAFs in inactive galaxies, and therefore under-estimate the fraction of local black hole mass accreted during ADAF phases. {The another one is that a different AGN population synthesis model for the newly measured XRB with \\textit{Swift} is adopted in the calculations in this work. Our present results are roughly consistent with that derived from observed Eddington ratio distributions in \\citet{h06}.}         ADAFs are very hot, and the temperature of the ions in ADAFs can be as high as $\\sim 10^{12}$K (see Fig. \\ref{structure_mdot0p001}), which implies that $\\gamma$-ray emission may be produced through the pion production processes in the proton-proton (p-p) collisions and subsequently decay of neutral pions \\citep{m97}. We calculate their contribution to the extragalactic $\\gamma$-ray background (EGRB), and find that less than 1\\% of the observed EGRB is contributed by the ADAFs in these faint sources if the massive black holes are spinning rapidly, while the contribution to the EGRB from ADAFs can be neglected if the black holes are non-rotating (see Figs. \\ref{fig_ele_g_delta0p1} and \\ref{fig_ele_g_delta0p01}). Our results are consistent with the previous works showing that about $\\sim 25\\%$ to $\\sim 100\\%$ of the EGRB can be attributed to the unresolved blazars \\citep*[e.g.,][]{p93,c95,s96,m00,c08,b09}."
    },
    "1005/1005.2196_arXiv.txt": {
        "abstract": "{ We study in detail (p)reheating after multi-field inflation models with a particular focus on N-flation.  We consider a variety of different couplings between the inflatons and the matter sector, including both quartic and trilinear interactions with a light scalar field.  We show that the presence of multiple oscillating inflatons makes parametric resonance inefficient in the case of the quartic interactions.  Moreover, perturbative processes do not permit a complete decay of the inflaton for this coupling.  In order to recover the hot big bang, we must instead consider trilinear couplings.  In this case we show that strong nonperturbative preheating is possible via multi-field tachyonic resonance.  In addition, late-time perturbative effects do permit a complete decay of the condensate. We also study the production of gauge fields for several prototype couplings, finding similar results to the trilinear scalar coupling.  During the course of our analysis we develop the mathematical theory of the quasi-periodic Mathieu equation, the multi-field generalization of the Floquet theory familiar from preheating after single field inflation.  We also elaborate on the theory of perturbative decays of a classical inflaton condensate, which is applicable in single-field models also. }  \\preprint{} ",
        "introduction": "\\label{sec:intro}  Reheating after inflation is a crucial requirement for any model.  Depending on how inflation ends and how the inflaton interacts, the transfer of energy from the inflaton sector into standard model (SM) degrees of freedom may proceed differently.  The detailed physics of (p)reheating is rather model dependent and previous studies have uncovered a wide array of different quantum field theory (QFT) phenomena, including perturbative \\cite{kofman25plus} and nonperturbative \\cite{KLS94,KLS97,frolov} production of both bosons and fermions \\cite{ferm}.  At the nonperturbative level a rich variety of preheating processes are possible, including parametric resonance \\cite{KLS94,KLS97,res}, tachyonic preheating \\cite{tac} and tachyonic resonance \\cite{tac_res}.  Examples have also been studied where the dynamics of reheating involves distinctively string theoretic processes \\cite{BBC,KY,BBHK} and cannot be captured by ordinary QFT.  (See also \\cite{stringy_reheating} for further studies of string theory reheating and \\cite{post_inf} for other aspects of the endpoint of string theory inflation.)  The requirement of successful reheating places a number of nontrivial constraints on inflationary model building.  At the most basic level, one must ensure that the inflaton field completely decays into radiation after inflation, in order to match smoothly onto the usual hot big bang evolution.  Moreover, one must verify that exotic relics (which might interfere with the successful predictions of big bang nucleosynthesis) are not over-produced.  It is therefore crucial to understand the decay channels of the inflaton in order to assess the viability of any inflation model.  Depending on the model, the process of preheating may also lead to various cosmological observables such as gravitational waves \\cite{gw} or nongaussianities \\cite{ng1,ng2}.  Previous studies of (p)reheating after inflation have focused largely on the decay of a small number of fields in the inflaton sector \\cite{KLS97}.  However, there are strong motivations to consider inflation models in which a huge number of fields play an important dynamical role.  One might expect such models to arise naturally in string theory, where realistic compactifications often contain exponentially large numbers of moduli and axion fields \\cite{bigN,Nflation,matrix_spectrum}.  Indeed, string theory appears to contain a rich landscape \\cite{Susskind} and, in general, stringy inflation might involve the motion of a huge number of fields on a very complicated scalar potential \\cite{multi_land,staggered,rand1,rand2,stream}.  Only in very special corners of this cosmic landscape would one expect inflation to take place along a single direction in field space.  Multi-field inflation is also desirable from a phenomenological perspective.  The ``assisted inflation'' mechanism (first discovered in \\cite{assisted} and generalized to a much broader class of models in \\cite{assisted2}) allows the collective motion of a large number of scalar fields $\\phi_i$ ($i = 1,\\cdots, N$) to drive inflation, even when each individual field could not support inflation on its own.  The essence of this mechanism is very simple. Suppose the fields are decoupled \\begin{equation} \\label{sumV} V(\\phi_j) = \\sum_{i=1}^N V_i(\\phi_i) \\, . \\end{equation} Then each inflaton feels the Hubble friction induced by the collective potential energies \\begin{equation} \\label{fried} 3 H^2 = \\frac{1}{M_p^2} \\sum_{i=1}^N \\left[ V_i(\\phi_i) + \\frac{\\dot{\\phi}_i^2}{2} \\right] \\cong \\frac{1}{M_p^2} \\sum_{i=1}^N  V_i(\\phi_i) \\end{equation} but only the restoring force coming from its own potential \\begin{equation} \\label{KG} \\ddot{\\phi}_i + 3 H \\dot{\\phi}_i = - V'_i(\\phi_i) \\, . \\end{equation} With sufficiently large $N$ the motion of each $\\phi_i$ will be over-damped, even if each $V_i(\\phi_i)$ is relatively steep, alleviating (to some extent) the fine-tuning problems associated with obtaining flat scalar potentials (see, for example, \\cite{flat_scalar_pot}). This mechanism permits a realization of chaotic inflation in a regime where each $\\phi_i$ undergoes a sub-Planckian displacement in field space, whereas the collective radial excitation $\\rho$ (defined by $\\rho^2 = \\sum_i \\phi_i^2$) traverses a super-Planckian distance.  Such a model goes a long way to ease concerns about the validity of the effective field theory description of large field inflation \\cite{lyth_bound}.  In \\cite{Nflation} a particularly appealing realization of assisted inflation was proposed: \\emph{N-flation}.  In this model one considers the dynamics of $N$ axion fields $\\phi_i$, each associated with a different broken shift symmetry $\\phi_i \\rightarrow \\phi_i + c_i$.  In the region relevant for inflation, the potential takes the simple form (\\ref{sumV}) with $V_i(\\phi_i) \\cong \\frac{m_i^2}{2}\\phi_i^2$.  The potential of N-flation has been argued to be radiatively stable \\cite{Nflation}.  (See also \\cite{Nflation2,Nflation3,Nflation4} for more details on the dynamics/phenomenology of this model.) N-flation is quite natural from the string theory perspective.  In oriented critical string theories the massless 2-form field $B_{\\mu\\nu}$ can wrap a huge number of independent 2-cycles when the theory is compactified to four dimensions.  Each such cycle leads to an axion at low energies.  Explicit examples are known where the number of low energy axions may be as large as $N = \\mathcal{O}(10^5)$ \\cite{bigN}.  % Using random matrix theory, it has been shown that the predictions of N-flation are largely independent of the details of the compactification \\cite{matrix_spectrum}.  See, however, \\cite{KSS} for complications associated with consistently embedding this framework into stabilized string theory vacua.  In this paper we study (p)reheating after multi-field inflation.  (See also \\cite{multi_preheat,multi_preheat2}.)  We focus on N-flation as a specific example, however, many of our results will apply to multi-field models more generally. We proceed phenomenologically, considering a variety of different couplings between the inflatons and matter.  For simplicity, we first consider reheating into a scalar field $\\chi$ through prototype interactions of the form \\begin{equation} \\mathcal{L}^{4-\\mathrm{leg}}_{\\mathrm{int}} = -\\sum_{i=1}^N \\frac{g_i^2}{2} \\phi_i^2 \\chi^2, \\hspace{5mm} \\mathcal{L}^{3-\\mathrm{leg}}_{\\mathrm{int}} = -\\sum_{i=1}^N \\frac{\\sigma_i}{2} \\phi_i \\chi^2 \\, . \\end{equation} These two kinds of couplings yield very different results for the decay of the inflaton, both in the perturbative and nonperturbative regimes.  In the case of 4-leg interactions, nonperturbative preheating processes are suppressed by the de-phasing of the N-flaton oscillations.  We show that reheating must therefore proceed via perturbative decays for this choice of coupling.  This is in contrast to the single-field case, were nonperturbative parametric resonance may occur.  On the other hand, we show that the 3-leg interactions \\emph{do} allow for violent nonperturbative particle production via tachyonic resonance.  The presence of a huge number of oscillating fields in the inflationary sector does \\emph{not} prohibit preheating via such couplings.  During the course of our analysis we develop the mathematical theory of the quasi-periodic Mathieu equation, the generalization of the usual Floquet theory to multi-field inflationary models.  We study quantitatively the structure of the stability/instability charts for multi-field inflation and comment on the dissolution of the stability bands in the presence of multiple incommensurate frequencies.  Having ruled out nonperturbative preheating effects in the case of the 4-leg interactions, we next study the perturbative QFT decays of the coherent oscillating N-flatons.  We find that the 4-leg interaction does \\emph{not} permit a complete decay; some relic density of inflaton particles always freezes out and the late-time universe is typically cold and devoid of life as we know it.  This result is completely analogous to the single-field case, which we re-consider.  Next, we study the perturbative decays of the coherent inflaton oscillations via 3-leg interactions, which will be relevant during the final stages of reheating and also when the inflaton oscillations have small amplitude.  Here we show that the inflaton condensate \\emph{is} permitted to decay completely and the late-time universe may be compatible with the usual hot big bang scenario.  Again, this result is completely analogous to the single-field case.  Finally, we also consider coupling the N-flaton particles to gauge fields.  We consider the following prototype interactions between the inflatons $\\phi_i$ and a $U(1)$ gauge field $A_\\mu$ with field strength $F_{\\mu\\nu} = \\partial_\\mu A_\\nu - \\partial_\\nu A_\\mu$: \\begin{equation} \\mathcal{L}^{\\mathrm{axion}}_{\\mathrm{int}} = -\\sum_{i=1}^N\\frac{\\phi_i}{4 M_i} F^{\\mu\\nu} \\tilde{F}_{\\mu\\nu}, \\hspace{5mm} \\mathcal{L}^{\\mathrm{moduli}}_{\\mathrm{int}} = -\\sum_{i=1}^N\\frac{\\phi_i}{4 M_i} F^{\\mu\\nu} F_{\\mu\\nu}  \\, . \\end{equation} Here $\\tilde{F}_{\\mu\\nu} = \\frac{1}{2} \\epsilon^{\\mu\\nu\\rho\\sigma} F_{\\rho\\sigma}$ is the dual field strength tensor, $\\epsilon^{\\mu\\nu\\rho\\sigma}$ is the anti-symmetric Levi-Civita symbol and the $M_i$'s are symmetry breaking scales. The former axion-type coupling to $F\\wedge F$ is directly relevant for N-flation.  In this case we find that the production of gauge fields is very similar to the 3-legs scalar interaction above; the inflaton is permitted to decay completely and nonperturbative preheating is possible via tachyonic resonance, depending on the model parameters $M_i$.  For completeness, we also consider the coupling $\\sum_i\\phi_i F^2/M_i$, which might be relevant for string theory models in which inflation is realized via the motion of a large number of moduli fields.  The organization of this paper is as follows. In section \\ref{sec:Nflation} we review the N-flation model and its string theoretic realizations.  In section \\ref{sec:preheating} we discuss the possibility of nonperturbative preheating after N-flation.  During the course of our analysis, we develop the mathematical theory of the quasi-periodic Mathieu equation, a multi-field generalization of the Floquet analysis that is familiar in the preheating literature. Next, we turn our attention to perturbative decays.  In section \\ref{sec:pert} we review the perturbative decays of the coherent oscillations of the inflaton in single field inflation and compare this to the more familiar case of decaying inflaton \\emph{particles}.  In section \\ref{sec:4legs} we discuss the perturbative decay of the N-flatons via 4-leg interactions, showing that the inflaton is not permitted to decay completely.  In section \\ref{sec:3legs} we consider instead perturbative 3-leg interactions showing that, this time, a complete decay is possible.  In section \\ref{sec:SM} we study the production of standard model gauge fields via both perturbative and nonperturbative processes after N-flation.  In section \\ref{sec:generalizations} we discuss some possible applications of our results outside of the context of N-flation and also study the production of gauge fields via ``modular'' couplings to the inflatons.  Finally, in section \\ref{sec:conclusions} we conclude and speculate on directions for future investigation. ",
        "conclusions": "\\label{sec:conclusions}   In this paper we have studied (p)reheating after inflation in models with a large number of inflaton fields. Such models are quite natural from the string theory perspective and may predict significant gravitational waves in a regime where the individual inflaton displacements are sub-Planckian, via the assisted inflation mechanism. We have focused on N-flation \\cite{Nflation} in particular, however, we expect that our results may apply more generally to a variety of multi-field inflation models.  In particular, our finding may have some relevance for multi-field inflation on the cosmic landscape \\cite{rand1,rand2,stream}.  We have proceeded phenomenologically, considering a variety of possible couplings between the inflaton fields and matter.  We studied the decay of the N-flatons into scalar field matter via both 4-leg and 3-leg interactions. In the former case, we found that reheating does not complete (similarly to single field inflation).  In the latter case, we found that reheating may complete in the perturbative regime, and also that strong nonperturbative tachyonic resonance is possible for certain parameters.  In the case of 3-leg interactions, the fluctuations of the preheat field display exponential growth in certain regions of phase space.  The structure of the stability/instability bands is quite rich and is related to the mathematical properties of  the quasi-periodic Mathieu equation. We have shown that the presence of multiple frequencies in the effective mass of the preheat field leads to interference effects and the dissolution of the stable regions.  %  We also studied the decay of the N-flatons to gauge fields.  In the original proposal \\cite{Nflation} the N-flatons are axions, therefore we considered a coupling of the form $\\sum_i \\phi_i F \\tilde{F}$.  We couple \\emph{all} of the N-flatons to a single $U(1)$ gauge field since the mass basis $\\phi_i$ and interaction basis $\\tilde{\\phi}_i$ are not, in general, aligned with one-another.  We have discussed how this interaction arises in the context of type IIB string theory vacua.  We found that (p)reheating in this case proceeds very similarly to the 3-legs scalar interaction.  Reheating can complete perturbatively and strong nonperturbative preheating effects are possible.  For completeness, we have also considered a coupling $\\sum_i\\phi_i F^2$ to gauge fields.  Such terms may arise in string theory models where assisted inflation is realized from the dynamics of moduli fields.  During the course of our analysis we have developed an analytical theory of the perturbative decays of coherent, classical inflaton oscillations.  This formalism applies much more generally than the N-flation models considered here.  We have used this method to compare the difference between the decay/annihilation of a classical, homogeneous inflaton condensate with the analogous process for inflaton \\emph{particles}.  We have seen that these two processes are essentially different.  In the case of a classical inflaton, the various decay channels of $\\phi$ contribute coherently to the amplitude.  Upon squaring this, we see that these decay channels can interfere quantum mechanically leading to oscillatory contributions to the effective Boltzmann equation.  These oscillations disappear upon averaging over a sufficiently long time interval and we recover the usual results for decay/annihilation of inflaton particles at rest.  There are a number of directions for future studies.  It would be interesting to consider a rigorous, stabilized stringy embedding of N-flation and identify the location of the standard model and the various decay channels of the inflaton \\cite{green}.  Using the results of this paper, one could assess the viability of such a construction and study the production of exotic relics, gravitational waves \\cite{gw}, nongaussianities \\cite{ng1,ng2}, etc. It would also be interesting to study post-inflationary dynamics in racetrack inflation \\cite{race1,race2}, or in a more general string landscape setting \\cite{rand1,rand2}.  There does not yet exist a complete theory of preheating after multi-field inflation.  The result of this paper should be considered as a step in the direction of a more comprehensive study.  As in the single field case, we expect that a rich spectrum of perturbative and nonperturbative QFT effects are possible \\cite{KLS97}, in addition to distinctively stringy processes \\cite{BBC,KY,BBHK}."
    },
    "1005/1005.2475_arXiv.txt": {
        "abstract": "We investigate the star formation and chemical evolution in the early universe by considering the merging history of the Galaxy in the $\\Lambda$ cold dark matter scenario according to the extended Press-Schechter theory. We give some possible constraints from comparisons with observation of extremely metal-poor (EMP) stars, made available by the recent large-scaled surveys and by the follow-up high-resolution spectroscopy. We demonstrate that (1) The hierarchical structure formation can explain the characteristics of the observed metallicity distribution function (MDF) including a break around $\\feoh=-4$. (2) A high mass initial mass function (IMF) of peak mass $\\sim 10 \\msun$ with the contribution of binaries, derived from the statistics of carbon enhanced EMP stars (Komiya et al. 2007, \\apj, 658, 367), predicts the frequency of low-mass survivors consistent with the number of EMP stars observed for $-4 \\lesssim \\feoh \\lesssim -2.5$. (3) The stars formed from primordial gas before the first supernova (SN) explosions in their host mini-halos are assigned to the hyper metal-poor (HMP) stars with $\\feoh \\sim -5$. (4) There is no indication of significant changes in the IMF and the binary contribution at metallicity  $-4\\gtrsim \\feoh \\gtrsim -2.5$, or even larger as long as the field stars of Galactic halo are concerned..  We further study the effects of the surface pollution through the accretion of interstellar matter (ISM) along the chemical and dynamical evolution of the Galaxy for low-mass population III and EMP survivors. Because of shallower potential of smaller halos, the accretion of ISM in the mini-halos in which these stars were born dominates the surface metal pollution. This can account for the surface iron abundances as observed for the HMP stars if the cooling and concentration of gas in their birth mini-halos is taken into account.  We also study the feedback effect from the very massive population III stars. The metal pre-pollution by pair-instability SNe (PISNe) is shown to be compatible with the observed lack of their nucleosynthetic signatures when some positive feedback on gas cooling works and changes IMF from being very massive to being high mass. ",
        "introduction": "In the early universe, the first stars should be formed from the primordial matter without metal elements \\citep[e.g.][and references therein]{Barkana01, Bromm04}. These stars formed without metals are referred to as Population III stars. Despite longstanding efforts, however, direct proof of their existence has been elusive. The first stars illuminate the universe and enrich surrounding matter with metals and the subsequent generations of stars with low metallicity should be formed under their influence. At the early stages of formation and evolution of galaxies, the feedback from stars is thought to play an important role. Since the era of the first stars and galaxy formation has occurred in the Dark Age of the universe, however, it is difficult to investigate them by direct observation with current instruments. One promising observational probe into this era is the extremely metal-poor (EMP) stars, which were stars formed in the early stages of the Galactic chemical evolution. The observations of EMP stars are expected to be connected with the first stars and Galaxy formation.  Over the past decade, thousands of very metal-poor stars below $\\feoh \\simeq -2$ have been identified in the Galactic halo by two large-scaled surveys (HK survey, Beers et al.\\ 1992: Hamburg/ESO [HES] survey, Christlieb et al.\\ 2001). The follow-up observations with high-dispersion spectroscopy (HDS) using large telescopes yield details of hundreds of these star's surface abundances \\citep[e.g., see the review by][]{Beers05a}. Thanks to these observations, now we are able to discuss statistical features of these stars \\citep[see, e.g., SAGA database;][]{Suda08, Suda09}. It is particularly worth noting that the most iron-deficient objects known to date are identified among the HES sample with much smaller metallicity than other EMP stars, i.e., with $\\feoh=-5.6$, $-5.4$ and $-4.8$ (HE1327-2326, Frebel et al.\\ 2005: HE0107-5240, Christlieb et al.\\ 2002: HE0557-4840, Norris et al.\\ 2007). We refer to these three stars as hyper metal-poor (HMP) stars, distinct from EMP stars, defined as stars with $\\feoh \\lesssim-2.5$.  The EMP stars display some features distinct from the Population~II stars and Pop.~III stars, as summarized in Section~\\ref{EMPs}. In previous studies \\citep[][, hereafter Paper I and Paper II, respectively]{Komiya07,Komiya09} we propose that they compose a peculiar stellar population formed with a high-mass initial mass function (IMF). We also argue that most of observed EMP stars were formed as the secondary members in binaries.  From the observations of EMP stars in the Galactic halo, the scarcity of stars below $\\feoh \\simeq -4$ has emerged and has been investigated from the perspective of chemical evolution \\citep[]{Prantzos03, Karlsson05}. Among $\\sim150$ stars with $\\feoh<-3$ with HDS observations, only three HMP stars are discovered below $\\feoh\\leqq-4.8$ and no stars are identified between $-4.8<\\feoh<-4$, although a few stars are found slightly below the $\\feoh=-4$. In Paper II, we propose a new scenario for this scarcity that the stars are divided into two groups according to whether they were formed before and after the parent mini-halos were polluted by their own first supernovae (SNe). HMP stars belong to the former group while the stars of the latter group constitute EMP stars. % We summarize this scenario in Section~\\ref{cutoffS}.  There have been investigations on the origins of the metal elements in the HMP stars. All of three HMP stars, known to date, show large carbon enhancement. \\citet{Suda04} propose a binary scenario to assert that HMP stars can be the survivors of the low-mass Pop.~III stars without pristine metal elements. In this scenario, their surface metal elements originate from afterbirth pollution by the accretion of interstellar matter (ISM) and by the wind accretion of envelope matter ejected by the binary companions. \\citet{Umeda03} and \\citet{Iwamoto05} propose that HMP stars are formed of the gas, enriched with metals by the ejecta of peculiar faint Pop.~III SN. \\citet{Limongi03} propose that a combination with two Pop.~III SNe can reproduce the abundance pattern peculiar to HMP stars. \\citet{Maynet06} suggest that light metal elements on HMP stars are brought by the stellar wind from rapidly rotating Pop.~III massive stars. For the binary scenario, \\citet{Nishimura09} investigate the AGB nucleosynthesis under the dearth of pristine metals to show that the abundance patterns of light metal elements from C through Al can be reproduced in terms of the accretion of envelope matter from AGB companion in binaries. In this case, the origin of the iron group elements remains to be discussed because binary mass transfer or stellar wind cannot provide these elements. As for the SN scenario, there remains a problem of forming the next generation of stars with an appropriate mixture of SN ejecta and ISM, and in particular, for faint SNe with the peculiar abundances, the stars have to be formed with the efficient use of a relatively small amount of SN ejecta. It is necessary, therefore, to discuss the origin of iron on the surface of HMP stars in the proper context of structure formation of the Galaxy.  From the theoretical viewpoint, many calculations of the formation of first stars in the early universe have been carried out \\citep[e.g.][]{Bromm99, Abel02, Omukai03, Tan04, O'Shea07, Yoshida08}. These studies suggest that one very massive star with $M \\gtrsim 100\\msun$ may be formed in the primordial mini-halos with mass $10^6 \\msun$, although typical masses or the mass range of the first stars are yet to be properly decided, especially in relation to the role of angular momentum in the fragmentation \\citep{Clark08,Machida08}. On the other hand, the evolution of Pop.~III and EMP stars have been investigated rather well, and in particular, it is shown that very massive stars with $M \\sim 200\\msun$ explode as pair-instability SNe (PISNe) with huge explosion energy and a large amount of metal yield \\citep{Umeda02, Heger02}. The effects of these energetic SNe on the subsequent star formation and galactic chemical evolution are not yet understood well, however \\citep[e.g.][]{Machida05, Kitayama05, Greif08}. The conditions for the formation of the first low-mass star are subject to controversy as well. Some studies propose that metal enrichment by massive stars accelerates the cooling of molecular clouds to lower the Jeans mass and to enable the low-mass star formation \\citep{BrommL03, Omukai05}. Other studies propose that ionization photons emitted from first stars accelerate the formation of ${\\rm H}_2$ molecules as main coolant in the primordial gas \\citep{Ricotti02, Yoshida08a}, which can trigger the first formation of low-mass stars. There remains the possibility that low-mass stars can be formed as the first objects \\citep{Nakamura01, Clark08}.  In this paper, we focus on the origin and nature of HMP and EMP stars and discuss the first stars, the early chemical evolution, and the star formation history of the Galaxy in the context of the hierarchical structure formation scenario. We have built the merger trees of the Galaxy using a semi-analytic method and follow the history of star formation and chemical enrichment along the trees \\citep[][hereafter, Paper III]{Komiya09L}. We register all the individual Pop.~III and EMP stars and trace their formation and evolution.  Firstly, we derive the MDFs and compare them with those of the observational counterparts in the Galactic halo. We adopt the high mass IMF, based on the comparisons of stellar models with observations, and take account of the binary contribution according to Paper I and Paper II. We also illustrate the chemical enrichment history in the hierarchical structure formation scenario and discuss the dependence of MDF on the assumptions about feedback from first stars and star formation rate in the early universe. In our study, we take account not only of the shape of the MDF but also of the total number of EMP and HMP stars.   Secondly, on the basis of this scenario, we study the effect of surface metal enrichment by accreting the ISM and attempt to predict the observational abundances of the first low-mass stars. Unlike the Pop.~I and Pop.~II stars, the ISM accretion can affect the surface metal abundance of Pop.~III and EMP stars because of their very low pristine metal abundances \\citep{Yoshii81, Iben83}. Furthermore, in small halos, in which Pop~III and EMP stars are born, the accretion rates of ISM gas are thought to have been much higher as compared with the those in the present Galactic halo because of much shallower gravitational potential. In order to estimate the ISM accretion rate and the change in the surface abundances of halo stars in the early universe, therefore, we should follow the merging history of halo as well as the star formation and chemical evolution. There are some previous studies of the ISM accretion \\citep{Iben83, Frebel09} onto Pop.~III stars but none of them take account of the dynamical and chemical evolution self-consistently. In this paper, we present the calculations of ISM accretion taking into account the merging history and the metal enrichment history of the Galaxy to explore the dependence of accretion rates on the dynamics of gas and stars. We demonstrate that the ISM accretion provides a reasonable interpretation of the observed iron abundance of HMP stars.  Thirdly we study possible metal pre-pollution by energetic SNe and draw observational constraints on the first stars and feedback from them.  % If PISNe or hypernovae (HNe) with the explosion energy and iron yield much larger than those of normal core collapse SNe \\citep{Heger02, Umeda02} occur in the low-mass halos, the blast wave may sweep out the gas from their host halos \\citep{Bromm03, Kitayama05, Machida05}. As a result, their metal ejecta is blown off from the host halos to pollute the inter-halo gas. We calculate the MDFs with such effect due to PISNe according to the hierarchical chemical evolution program and compare them with the observed HMP and EMP stars.  A semi-analytic approach to hierarchical structure formation has been adopted by \\citet{Tumlinson06} and by \\citet{Salvadori06}, who have constructed the merger trees of the Galaxy and investigated inhomogeneous Galactic chemical evolution models. \\citet{Tumlinson07L} considered an IMF dependent on the temperature of the cosmic microwave background, and discussed the carbon-rich EMP stars also by taking into account the contribution of binaries. However, he did not relate hierarchical nature to the scarcity of the HMP stars. We address the difference between the HMP and EMP stars in the context of hierarchical galaxy formation. Additionally, we discuss the total number of EMP stars as well as shape of the MDF. Furthermore, none of these studies is capable of explaining the origin of the extraordinarily low iron abundance of HMP stars. We examine two possible origins, i.e., the surface pollution by ISM accretion and the pre-pollution of intra-Galactic matter (IGM) by halo blown off, for the small amount of iron in HMP stars.  The paper is organized as follows. We review the observations and elaborate our scenario in the next section and describe the computational method of star formation and evolution of Galaxy, in Section~\\ref{numericS}. The following three sections are devoted to presenting the results of our computations for the halo formation models, the surface pollution due to the ISM accretion, and the effects of pre-pollution of mini-halos by energetic PISNe, respectively. We discuss the formation of the first stars and HMP stars in Section~\\ref{discussion} and a summary of the conclusion follows in Section~\\ref{summary}. ",
        "conclusions": "\\label{summary}  We have studied star formation and chemical evolution in the early universe from the perspective of the hierarchical structure formation of the Galaxy by constructing the merging trees on the $\\Lambda$ cold dark matter scenario according to the extended Press-Schechter theory. In particular, we investigate the formation history of EMP and HMP stars and the origin of iron in HMP stars. We demonstrate that the hierarchical chemical evolution can reproduce the characteristics of MDF and the frequency of EMP stars in the Galactic halo by applying the high-mass initial mass function and the contribution of binary, derived for EMP stars in our previous work \\citep{Komiya07,Komiya09}.  Our main conclusions are summarized as follows.  \\begin{enumerate} \\item In the hierarchical structure formation scenario, the metal-poor stars are divided into two groups, according to whether they are born before or after the first pollution, i.e., the first SNe pollution of the gas in each mini-halo. The stars of the former group are made of gas with the pristine abundances of mini-halos, and the stars of the latter group are enriched with metals, produced by their own SNe of each halo. The latter stars well reproduce the observed metallicity distribution of EMP stars of $\\feoh \\gtrsim -4$ and the observational counterparts of the former stars are found in the HMP stars of $\\feoh < -4$.  \\item While star formation efficiency has few effects except for the stars of the lowest metallicity, the initial mass function directly influences the frequencies of low-mass stars, survived to date for given metallicity. For EMP stars of $\\feoh \\gtrsim -4$, the high-mass IMF of medium mass $\\sim 10 \\msun$ with the contribution of binaries, derived on the basis of the statistics of CEMP and EMP stars by \\citet{Komiya07, Komiya09}, gives the total number of EMP stars, compatible with the observations by HES surveys. In particular, the low-mass stars observed in the Galactic halo today are mostly (more than 90\\%) born as the secondary members of binaries. On the other hand, the low-mass IMF such as derived for the Galactic spheroid component gives rise to the overabundance of EMP survivors by more than two orders of magnitude.  \\item For stars born before the first pollution, the scarcity of HMP stars suggests that the star formation, especially of low-mass survivors, is less efficient as compared with that for the EMP population of $-4 \\lesssim \\feoh \\lesssim -2.5$. The overproduction of low-mass survivors still persists by a factor of $\\sim 10$ even if we take into account the suppression of the star formation efficiency due to the local radiation feedback of the first massive stars. This is indicative of higher-mass IMF and/or less efficient binary formation in the gas of the lowest metallicity than for the EMP populations.  \\item As for higher metallicity of $\\feoh \\gtrsim -2.5$, there is no indication of significant changes in the IMF, in particular into a low-mass IMF as long as the fields stars in the Galactic halo are concerned. From the comparison with the MDF derived by \\citet{Ryan91} for kinematically selected samples, however, the star formation efficiency declines for $\\feoh \\gtrsim -1.8$. This occurs near the redshift $z \\sim 10$ in our fiducial model and may be related to the formation of the Galactic bulge and/or the Galactic discs.   \\item For the first low-mass stars, we have evaluated the effect of ISM accretion after birth by taking into account the dynamical and chemical evolution of the Galaxy. The accretion in the mini-halos, in which Pop.~III stars are formed, dominates the surface pollution, giving a much larger amount of accreted gas and metals than in the evolved massive halos. If we take into account the cooling of gas and resultant concentration of gas and stars into the central region of their birth halos, the ISM accretion can give rise to the surface pollution sufficient to explain the iron abundances observed for HMP stars.  \\item We have investigated the feedback effect from primordial stars, born prior to EMP stars, and the pre-pollution of intergalactic gas by SNe in the first collapsed halos. Primordial supernovae pollute intergalactic matter to $\\feoh=-3.2 - -3.8$ in our models and some EMP survivors with such metallicity can be formed with pre-pollution before the first pollution. Metal ejecta from Type II SNe that occurred in the mini-halos with positive feedback also pollute intergalactic matter. In the models with significant feedback effects, the ejecta from PISN with a characteristic abundance pattern should blend and can be consistent with observational absence of abundance pattern peculiar to PISNe. Detection of 3 HMP stars indicates that some positive feedback enables the low mass star formation below $\\feoh<-4$. \\end{enumerate}  From the binary scenario, we have revealed that there are two possible origins of iron for HMP stars, the surface pollution of accreted IMF in the mini-halos and the pre-pollution of intergalactic gas with metals that energetic SNe have blown off the host halos. The polluted Pop.~III stars with ISM accretion and the pre-polluted HMP stars differ only in the presence of pristine iron in their interior. In order to distinguish these two possibilities, further investigations are necessary to reveal the effects that the pristine metals of such small abundances as corresponding to the iron abundance observed for HMP stars have on the nucleosynthesis and evolution of low-mass and intermediate-mass stars \\citep[e.g., see][]{Suda09b}.  In this paper, we use the simplified assumption about the star formation and deal only with iron. In particular, we assume the instantaneous and homogeneous mixing of SN ejecta in the halos and in the IGM. In actuality, however, we should take into account the dependence of the expansion of shells on the explosion energy of SNe and host halo mass, and the finite timescale of mixing \\citep{Greif08, Tornatore07}. It little affects the present results as long as the host halos are not disrupted by SNe, however, since the SNe shells involve the circumstellar matter comparable to the smallest masses in the mini-halos that can nurture the star formations, except that it may reduce the number of stars with the smallest metallicity of EMP stars, which are formed in the halos of larger masses that collapse later. From our homogeneous mixing model, we cannot compute detailed MDF of stars formed before the first pollution and additional investigation is required to distinguish the origin of positive feedback. Further elaborations are necessary for the interactions between the star formation and SNe and the binary formation under the metal-deficient condition. We may expect to gain a better understanding of EMP stars and the evolution of the Galaxy from the large-scaled, deep surveys of halo stars now in progress such as SDSS/SEGUE and LAMOST, as well as from the detailed abundance studies with large telescopes.  This work is supported partly by JSPS Grand-in-aid for Scientific Research (18104003, 18072001, 19740098). A part of results are reported in AIP~Conference~990 \"First Stars\", AIP~Conference~1016 \"The Origin of Matter and Evolution of Galaxy\", and IAU~Symposium~255 \"Low-Metallicity Star Formation: from First Stars to Dwarf Galaxies\"."
    },
    "1005/1005.2182_arXiv.txt": {
        "abstract": "{Over the past few years several studies have provided estimates of the SFR (star--formation rate) or the total infrared luminosity from just one infrared band. However these relations are generally derived for entire galaxies, which are known to contain a large scale diffuse emission that is not necessarily related to the latest star--formation episode.}{We provide new relations to estimate the SFR from resolved star--forming regions at 100~$\\mu$m and 160~$\\mu$m.}{We select individual star--forming regions in the nearby (840~kpc) galaxy M33. We estimate the SFR combining the emission in H$\\alpha$ and at 24~$\\mu$m to calibrate the emission at 100~$\\mu$m and 160~$\\mu$m as SFR estimators, as mapped with PACS/{\\em Herschel}. The data are obtained in the framework of the HERM33ES open time key project.}{There is less emission in the HII regions at 160~$\\mu$m than at 100~$\\mu$m. Over a dynamic range of almost 2~dex in $\\Sigma_{SFR}$ we find that the 100~$\\mu$m emission is a nearly linear estimator of the SFR, whereas that at 160~$\\mu$m is slightly superlinear.}{The behaviour of individual star--forming regions is surprisingly similar to that of entire galaxies. At high $\\Sigma_{SFR}$, star formation drives the dust temperature, whereas uncertainties and variations in radiation--transfer and dust--heated processes dominate at low $\\Sigma_{SFR}$. Detailed modelling of both galaxies and individual star forming regions will be needed to interpret similarities and differences between the two and assess the fraction of diffuse emission in galaxies.} ",
        "introduction": "Star formation is one of the main drivers of galaxy formation and evolution and as such the accuracy of the SFR (star--formation rate) determination is of great importance for deriving the cosmic history of galaxies. Along with the UV and the H$\\alpha$, the total infrared luminosity is widely used to estimate the SFR. To properly quantify the infrared luminosity a good sampling of the infrared SED (spectral energy distribution) is needed \\citep{dale2002a,draine2007a}. Over the past few years, many authors have shown that the TIR (total infrared) luminosity, and by extension the SFR \\citep{kennicutt1998a}, can also be evaluated from monochromatic emission measures \\citep[][Y. Li et al. 2010, in prep.]{takeuchi2005a,calzetti2007a,rieke2009a,calzetti2010a,boquien2010a}. However, most relations between the far--infrared luminosity and the SFR are established for entire galaxies. At shorter wavelengths, some relations have been derived from individual star--forming regions \\citep{calzetti2005a,perez2006a,calzetti2007a,relano2007a}. As new deep surveys will become available at wavelengths where most of the energy is emitted, it is prudent to try to understand the relation between the integral emission of a galaxy and that of the individual star forming regions. The physical conditions, such as temperature, abundance and emissivity, of the infrared--emitting dust can vary widely in a galaxy, and so does the significant contribution from evolved stars \\citep{lonsdale1987a,sauvage1992a,buat1996a}. Therefore, any scaling relation established from the emission of entire galaxies may not be appropriate when applied to resolved star--forming regions in these same galaxies.  The limited resolution of far--infrared instruments onboard IRAS (Infrared Astronomical Observatory), ISO (Infrared Space Observatory) or even {\\em Spitzer} beyond 60~$\\mu$m made the study of individual star--forming regions within galaxies difficult. However, the recently launched {\\em Herschel} Space Observatory \\citep{pilbratt2010a} with its unprecedented resolution provides the first opportunity to study the spatially resolved far--infrared dust emission in exquisite detail. Such a fine resolution is of the utmost importance to study the emission of star--forming regions located in nearby galaxies in order to provide insights into the fundamental properties of the dust and to quantify the SFR.  With an inclination of 56$^{\\circ}$ \\citep{regan1994a} and a distance of only 840~kpc \\citep{freedman1991a}, M33 is one of the closest spiral galaxies. It has been imaged by {\\em Herschel} in the context of the HERM33ES Open Time Key Project \\citep{kramer2010a}, providing one of the finest views {\\em Herschel} will ever provide of a spiral galaxy from 100~$\\mu$m to 500~$\\mu$m \\citep{kramer2010a,braine2010a,verley2010b}. ",
        "conclusions": "\\label{sec:conclusion} We used the high--resolution {\\em Herschel} 100~$\\mu$m and 160~$\\mu$m observations of a nearby star--forming galaxy, M33. We combined {\\em Herschel} PACS data with {\\em Spitzer} MIPS 24~$\\mu$m and ground--based H$\\alpha$ to provide new calibrations of the 100~$\\mu$m and 160~$\\mu$m to estimate $\\Sigma_{SFR}$ from individual star--forming regions. For the selected star--forming regions in M33, the 100~$\\mu$m luminosity is a linear SFR estimator over a factor 100 in surface brightness, whereas the 160 micron luminosity is slightly superlinear. It appears that individual star forming regions exhibit a similar behaviour as entire galaxies taken from the LVL, SINGS, starburst galaxies from \\cite{engelbracht2008a} samples when estimating $\\Sigma_{SFR}$ from the 100~$\\mu$m and 160~$\\mu$m bands emission. In a similar fashion, the dust temperature -- as measured by the ratio of the 160~$\\mu$m to 100~$\\mu$m emission -- increases as a function of $\\Sigma_{SFR}$, suggesting that at high $\\Sigma_{SFR}$ the star formation drives the trend for both systems, while at lower $\\Sigma_{SFR}$ uncertainties and variations of conditions in the radiation--transfer and dust--heating processes contribute to the scatter. In other words, the fairly wide dust temperature distribution at low $\\Sigma_{SFR }$ becomes increasingly biased towards higher temperatures at higher $\\Sigma_{SFR}$ in both HII regions and entire galaxies."
    },
    "1005/1005.4137_arXiv.txt": {
        "abstract": "Most models of solar eruptions assume that coronal field lines are anchored in the dense photosphere and thus the photospheric magnetic fields would not have rapid, irreversible changes associated with eruptions resulted from the coronal magnetic reconnection. Motivated by the recent work of Hudson, Fisher \\& Welsch (2008) on quantitatively evaluating the back reaction due to energy release from the coronal fields, in this Letter we synthesize our previous studies and present analysis of new events about flare-related changes of photospheric magnetic fields. For the 11 X-class flares where vector magnetograms are available, we always find an increase of transverse field at the polarity inversion line (PIL) although only 4 events had measurements with 1 minute temporal resolution. We also discuss 18 events with 1 minute cadence line-of-sight magnetogram observation, which all show prominent changes of magnetic flux contained in the flaring $\\delta$ spot region. Except in one case, the observed limb-ward flux increases while disk-ward flux decreases rapidly and irreversibly after flares. This observational evidence provides support, either directly or indirectly, for the theory and prediction of Hudson, Fisher \\& Welsch that the photospheric magnetic fields must respond to coronal field restructuring and turn to a more horizontal state near the PIL after eruptions. ",
        "introduction": "Solar eruptions, including flares, filament eruptions, and coronal mass ejections (CMEs) have been understood as the result of magnetic reconnection in the solar corona (e.g., Kopp \\& Pneuman 1976; Antiochos et al. 1999). Although surface magnetic field evolution (such as new flux emergence and shear motion) play important roles in building energy and triggering eruption, most models of flares and CMEs have the implication that photospheric magnetic fields do not have rapid, irreversible changes associated with the eruptions. The key reason behind this assumption is that the solar surface, where the coronal magnetic fields are anchored, has much higher density and gas pressure than the corona. Recently, we note the work by Hudson, Fisher \\& Welsch (2008, hereafter HFW08), who quantitatively assessed the back reaction on the solar surface and interior resulting from the coronal field evolution required to release energy, and made the prediction that after flares, the photospheric magnetic fields become more horizontal. Their analysis is based on the simple principle that any change of magnetic field energy must lead to a corresponding change in magnetic pressure. This is one of the very few models that specifically predict that flares can be accompanied by rapid and irreversible changes of photospheric magnetic fields. Over a decade ago, Melrose (1997) used the concept of reconnection between two current-carrying systems to provide explanation for the enhancement of magnetic shear at the flaring magnetic polarity inversion line (PIL), which is sometimes observed (see below). Perhaps this is in the same line as the tether-cutting model for sigmoids, which was elaborated by Moore et al. (2001) and involves a two-stage reconnection processes. At the eruption onset, the near-surface reconnection between the two sigmoid elbows produces a low-lying shorter loop across the PIL and a larger twisted flux rope connecting the two far ends of the sigmoid. The second stage reconnection occurs when the large-scale loop cuts through the arcade fields causing erupting plasmoid to become CME and precipitation of electrons to form flare ribbons. If scrutinizing the magnetic topology close to the surface, one would find a permanent change of magnetic fields that conforms to the scenario of HFW08: the magnetic fields turn more horizontal near the flaring PIL due to the newly formed short loops there.  \\begin{deluxetable*}{lccccc} \\tablewidth{0pt} \\tablecaption{Events with Evidences of Enhancement of Transverse Magnetic Field ($B_t$) at the PIL} \\tablehead{\\colhead{Event Date}  & \\colhead{{\\it GOES} Level} & \\colhead{Data Source} & \\colhead{Data Cadence} & \\colhead{Main Findings} & \\colhead{References}} \\startdata (1) 1990 Aug 27   & X3.0  & BBSO  & 10 min & magnetic shear increase & 1 \\\\ (2) 1991 Mar 22   & X9.0  & BBSO  & 10 min & magnetic shear increase & 2 \\\\ (3) 1991 Jun 9   & X10.0  & BBSO/HSO  & 5 hr & magnetic shear increase & 2 \\\\ (4) 2000 Jul 14   & X5.7  & HSO  & 3 hr & $B_t$ and electric current increase & 3 \\\\ (5) 2001 Aug 25   & X5.3  & BBSO & 1 min & $B_t$ and magnetic shear increase & 4 \\\\ (6) 2001 Oct 19   & X1.6  & BBSO  & 1 min & $B_t$ and magnetic shear increase & 4 \\\\ (7) 2002 Jul 26   & M8.7  & BBSO & 1 min & $B_t$ increase & 5 \\\\ (8) 2003 Oct 29   & X10  & MSFC & 3 hr & $B_t$ and magnetic shear increase & 6 \\\\ (9) 2005 Jan 15   & X2.6  & BBSO & 1 min & $B_t$ and inclination angle increase & 7 \\\\ (10) 2005 Sep 13   & X1.5  & BBSO & 1 min & $B_t$ increase & 8 \\\\ (11) 2006 Dec 13   & X3.4  & Hinode  & 8 hr & $B_t$ and magnetic shear increase & 9 \\enddata \\tablecomments{References: (1) Wang 1992; (2) Wang et al. 1994; (3) Wang et al. 2005; (4) Wang et al. 2002; (5) Wang et al. 2004a; (6) Liu et al. 2005; (7) Li 2010; (8) Wang et al. 2007; (9) Jing et al. 2008} \\end{deluxetable*}  On the observational side, earlier studies were inconclusive on the flare-related changes of photospheric magnetic field topology. Wang (1992) and Wang et al. (1994) showed impulsive changes of vector fields after flares including some unexpected patterns such as increase of magnetic shear along the PIL, while mixed results were also reported (Ambastha et al. 1993; Hagyard et al. 1999; Chen et al. 1994, Li et al. 2000a, 2000b). It is not until recently that rapid and permanent changes of photospheric magnetic fields, mainly the line-of-sight component, are observed to be consistently appear in major flares and considered as indicative of flare energy release (Kosovichev \\& Zharkova 2001; Sudol \\& Harvey 2005). In particular, a number of papers of Big Bear Solar Observatory (BBSO)/New Jersey Institute of Technology group have been devoted to the finding of sudden unbalanced magnetic flux change (Spirock et al. 2002; Wang et al. 2002; Yurchyshyn et al. 2004; Wang et al. 2004a; Wang 2006) and a new phenomenon of sunspot white-light structure change (Wang et al. 2004b; Deng et al. 2005; Liu et al. 2005; Chen et al. 2006) associated with flares. By evaluating the mean relative motions between two magnetic polarities of five flaring $\\delta$ spots, Wang (2006) revealed a sudden release of the overall magnetic shear and found a correspondence between converging/diverging motion and increase/decrease of magnetic gradient at the PIL, which suggests magnetic reconnection at or close to the photosphere. Liu et al. (2005) discussed outer penumbral decay and central penumbral darkening of $\\delta$ spot seen in seven X-class flares, and proposed a reconnection picture where the active region field collapses inward after flares as signified by HFW08.  Although studying the change of magnetic field has promise to advance the understanding of flare energy release process, it is notable that the long-accumulated observational evidence have not yet converged regarding how the observed changes of photospheric magnetic fields could be understood in the context of models of coronal magnetic field restructuring. The objective of this Letter is hence to examine the observations in a systematic fashion and compare them quantitatively with the prediction of HFW08. Meanwhile, several conflicting concepts in our earlier papers will be reconciled with new physical understanding. We anticipate that these will guide more efficient studies to analyze vector magnetograph data of {\\it Helioseismology and Magnetic Investigation} ({\\it HMI}) on board {\\it Solar Dynamic Observatory} ({\\it SDO}) in the near future.  \\begin{figure} \\epsscale{1.15} \\plotone{f1.eps} \\caption{Time profiles of transverse field, inclination angle, and Lorentz force per unit area within a white circled region (in $a$) at the PIL (black line in $a$) for the 2002 July 26 M8.7 flare (see an H$\\alpha$ image in $b$), calculated using BBSO vector magnetograms. In {\\it GOES} 10 soft X-ray flux (dashed line), the flare started at 20:51~UT, peaked at 21:12~UT, and ended at 21:29~UT.} \\end{figure}  \\begin{deluxetable*}{lcccccc} \\tablewidth{0pt} \\tabletypesize{\\scriptsize} \\tablecaption{Events with Evidences of Rapid Changes of Active Region Line-of-Sight Magnetic Fluxes \\label{table 2} } \\tablehead{\\colhead{Event}  & \\colhead{NOAA AR} & \\colhead{Location (deg)} & \\colhead{{\\it GOES} Level} & \\colhead{Limbward Flux (10$^{20}$ Mx)} & \\colhead{Diskward Flux (10$^{20}$ Mx)} & \\colhead{References}} \\startdata (1) 1991 Mar 22 & 06555  & E20 S23  & X9.0  & $+1.0$ & 0 & 1 \\\\ (2) 2001 Apr 2 & 09393  & W64 N19  & X20.0  & $+6.0$ & $-1.5$ & 1, 4 \\\\ (3) 2001 Apr 6 & 09415  & E30 S20  & X5.6  & $+4.0$ & 0.0 & 1 \\\\ (4) 2001 Apr 9 & 09415  & W04 S21  & M8.0  & $+2.7$ & 0 & 2 \\\\ (5) 2001 Aug 25 & 09591  & E34 S17  & X5.3  & $+1.8$ & $-0.8$ & 1 \\\\ (6) 2001 Sep 24 & 09632  & E18 S18  & X2.6  & $+10$ & $-10$ & 5 \\\\ (7) 2001 Oct 19 & 09661  & W29 N15  & X1.6  & $+3.0$ & $-0.4$ & 1 \\\\ (8) 2001 Oct 22 & 09672  & E16 S18\t&  X1.2\t& $+11$\t& $-2$ & 1\\\\ (9) 2001 Oct 25 & 09672 & W26 S19 & X1.3 &\t$-3$ & $+5$ & 5\\\\ (10) 2002 Jul 23 & 10039 & E54 S12 &\tX4.8 &\t$+0.15$\t& $-0.8$ & 3\\\\ (11) 2003 Mar 18 & 10314 & W53 S16\t& X1.5\t& $+5.0$ & 0 &5 \\\\ (12) 2003 Oct 19 & 10484 & E54 N05\t& X1.1\t& $+7$ & $-2$ & 1\\\\ (13) 2003 Oct 26 & 10484 & N04 W41 & X1.2 & 0 & 0  & 5\\\\ (14) 2003 Oct 28 & 10486 & E04 S16 & X17.2 & $+0.75$ & $-0.4$ & 2\\\\ (15) 2003 Oct 29 & 10486 & W10 S17 &\tX10.0\t& $+0.65$ & $-0.33$ & 2\\\\ (16) 2004 Jul 16 & 10649 & S10 E26 & X3.6 & $+2.5$ & $-10.0$ & 5\\\\ (17) 2004 Nov 7 & 10696 & N09 W08 & X2.0 & $+5.0$ & $-3.0$ &5 \\\\ (18) 2005 Sep 13 & 10808 & E17 S11 &\tX1.5 & $+0.75$ & $-1.2$ & 2\\ \\enddata \\tablecomments{References: (1) Wang et al. 2002; (2) Wang 2006; (3) Yurchyshyn et al. 2004; (4) Spirock et al. 2002; (5) This paper} \\end{deluxetable*} ",
        "conclusions": "Synthesizing the research of flare-related rapid and irreversible changes in both vector and line-of-sight magnetic fields, we have revealed that photospheric fields apparently respond to the back reaction of coronal fields due to flare energy release, in a pattern that magnetic fields near the PIL become more horizontal, which strongly evidences the conjecture of HFW08. The change-over time lies between $\\sim$10 minutes to 1 hour, and all the changes are co-temporal with the flare initiation. These should provide insight into the upcoming {\\it HMI} data analysis towards a more complete physical understanding of the impacts of flares on the solar atmosphere following coronal field restructuring. We discuss a few points related to observations and modeling work.  1. It is obviously reasonable that larger flares tend to produce more prominent field structure change (e.g., Chen et al. 2006). Using unprecedented {\\it HMI} measurements, we anticipate finding a lower threshold of flare magnitude at which the magnetic field changes are still detectable.  2. A related aspect is that Hudson (2000) postulated that the energy conversion process during coronal transients would involve a magnetic implosion, which signifies that the flare ribbons and loops may contract initially before expansion. Although multi-wavelength signatures of such contraction have been observed recently (e.g., Liu \\& Wang 2009 and references therein), it is still unclear how to link it to the changing of field line orientation from more vertical to more horizontal configuration.  3. As our previous studies have shown, it is most likely that the strengthening of transverse fields near the PIL has a causal relationship to the enhancement of central penumbral structure of $\\delta$ spots (e.g., Wang et al. 2004b; Liu et al. 2005). However, the penumbral decay in the outer boundary region is not reflected in the model of HFW08. We speculate that it is due to the peripheral field lines changing to more vertical state when the central region pressure is released after flares. That is to say, the surrounding fields might be subsequently pushed inwards to fill the void.  4. HFW08 introduced the concept of a ``jerk'' (i.e., downward push in the vertical direction due to the change of Lorentz force), which may account for the launch of seismic waves seen in some impulsive flares (Kosovichev \\& Zharkova 1998). This connection has other supporting evidence (e.g. Mart\\'inez-Oliveros \\& Donea 2009). A pertinent phenomenon is that a sunspot can have rapid flare-induced motion along the solar surface, which seems to be sufficiently driven by the changes of the horizontal components of Lorentz force (Anwar et al. 1994). Liu et al. (2010, in preparation) presented five new events in an effort to associate the kinetic energy of sunspot lateral motion with that of seismic waves, while we must note that the detailed mechanism involved has never been researched before. In addition, although the events with seismic waves always have associated rapid changes of surface magnetic fields, it seems that this conclusion cannot be reversed as suggested by several cases under study. The new research area will shed more light on understanding the linkage between the sub-surface/surface magnetic activities and coronal eruptions."
    },
    "1005/1005.4071_arXiv.txt": {
        "abstract": "The \\hi 21~cm transition line is expected to be an important probe into the cosmic dark ages and epoch of reionization. Foreground source removal is one of the principal challenges for the detection of this signal. This paper investigates the extragalactic point source contamination and how accurately bright sources ($\\gtrsim 1$ ~Jy) must be removed in order to detect 21~cm emission with upcoming radio telescopes such as the Murchison Widefield Array (MWA). We consider the residual contamination in 21~cm maps and power spectra due to position errors in the sky-model for bright sources, as well as frequency independent calibration errors. We find that a source position accuracy of $0.1$~arcsec will suffice for detection of the \\hi power spectrum. For calibration errors, $ 0.05~\\%$ accuracy in antenna gain amplitude is required in order to detect the cosmic signal. Both sources of subtraction error produce residuals that are localized to small angular scales, $\\kperp \\gtrsim 0.05 $~Mpc$^{-1}$, in the two-dimensional power spectrum. ",
        "introduction": "The cosmological epoch of reionization (EoR) is a key milestone in the history of structure formation, marking the transition from a fully neutral to a highly ionized intergalactic medium (IGM) due to the ultra-violet and X-ray radiation of early stars, galaxies, and black holes.  Recent observations of the Gunn-Peterson effect, i.e., Ly$\\alpha$ absorption by the neutral IGM, toward the most distant quasars ($z \\sim 6$), and the large scale polarization of the CMB, corresponding to Thompson scattering during reionization, have set the first constraints on the reionization process. These results suggest significant variance in both space and time, starting perhaps as far back as $z \\sim 11$ (\\citet{komatsu10}; WMAP 7 year data) and extending to $z \\sim 6$ \\citep{fan06b}. Previous WMAP five year data indicates the 5~$\\sigma$ detection of the E-mode of polarization which rules out any instantaneous reionization at $z \\sim 6$ at 3.5~$\\sigma$ level. In case of the Gunn-Peterson effect, the IGM becomes optically thick to Ly$\\alpha$ absorption for a neutral fraction as small as $\\sim 10^{-3}$. In order to overcome these limitations, it has been widely recognized that mapping the red-shifted \\hi 21~cm line has great potential for direct studies of the neutral IGM during reionization \\citep{barkana01, fan06b, furlanetto06, morales09}.  There are number of upcoming low-frequency arrays with key science goals to detect the \\hi 21~cm signal from the EoR. This includes the Murchison Widefield Array [MWA] \\citep{mitchell08, lonsdale09, bowman06}, Precision Array to Probe Epoch of Reionization [PAPER] \\citep{backer07, parsons09a}, Low Frequency Array [LOFAR] \\citep{harker10, jelic08, panos09} and Giant Meterwave Radio Telescope [GMRT] \\citep{pen09}.  One of the major challenges for all of these upcoming arrays will be the removal of the continuum foreground sources in order to detect the faint \\hi signal from the EoR.  A variety of continuum foregrounds complicate redshifted 21~cm measurements of the EoR \\citep{shaver99}. Diffuse Galactic synchrotron emission dominates the low-frequency radio sky and is approximately four orders of magnitude brighter than the $\\sim10$~mK 21~cm signal at the frequency relevant to reionization ($\\nu\\approx150$~MHz).  In addition, Galactic and extragalactic free-free emission contribute additional flux to the diffuse foreground.  Radio point sources from AGN, radio galaxies, and local Galactic sources are numerous and particularly challenging.  The brightest of these sources have fluxes well above $S>1$~Jy and are seven or eight orders of magnitude above the EoR signal in low-frequency radio maps.  The distribution of point sources also extends to very faint levels such that the brightness temperature due to confused sources in upcoming arrays will be $\\sim10$~K, or three orders of magnitude brighter the than the 21~cm background.  In this paper we discuss how the radio interferometric imaging techniques are going to affect the foreground source modeling and subsequent removal from the data-set in order to search for the EoR signal. Recently, there has been extensive research on foreground source modeling at these low frequencies \\citep{dimatteo02, jelic08, rajat09}. Similar effort has also been made in exploring different techniques to remove the foregrounds from the EoR data-set by \\citet{morales06a,morales06b, bowman09, liu09, parsons09b, gleser08, harker09a, harker09b}. Since attempting to observe a signal below the confusion limit of foreground sources is a novel aspect of 21~cm experiments, most of these works primarily focus on the removal of faint and confused sources that fall below a specified cutoff flux limit, $S_{cut}$ ($\\approx 1$~Jy). They do not consider the foreground sources brighter than $S_{cut}$ and how accurately they need to be removed.  Indeed, most of these analysis implicitly assume that the bright foreground sources above $S_{cut}$ have been removed perfectly. But in reality imperfect instrument calibration or any errors in the subtracted foreground model will introduce artifacts and leave residual contamination in the data after bright source removal, even by traditional techniques such as ``peeling''.  These residuals may interfere with either the subsequent faint source subtraction or the ultimate detection and characterization of the redshifted 21~cm signal.  In \\citet{abhi09} we dealt with the bright point sources above $S_{cut}$ and the limitations that will be caused due to imperfect removal of such sources in the image plane. In this paper we extend the initial analysis in order to estimate the residual contamination in the power spectral domain of improper bright source subtraction. The objective of this paper is to demonstrate how the accuracy in the foreground removal affects the detection of \\hi 21~cm power spectra with the MWA.  In Section 2, we discuss our choice of sky model and outline the simulation parameters, including the array specifications and data reduction procedure, and describe the two categories of corruption terms that we will consider: source position errors and residual calibration errors. The results obtained for the residual angular power spectrum, spherically averaged three-dimensional power spectrum, and two-dimensional power spectrum are presented in Section~3. Finally, in the last section we summarize the implications of the results from our simulations. ",
        "conclusions": "With the results from Section~\\ref{s_results}, we can address our three key questions for bright source subtraction residuals in the power spectral domain.  First, we find that the level of the source subtraction accuracy required for power spectral detection of the 21~cm signal is roughly comparable to the accuracy required for direct imaging of the \\hi signal \\citep{abhi09}.  The power spectrum tolerance does suffer, however, compared to the direct imaging case in one regard.  For direct imaging, it is possible to find areas in the final image map that are far from bright sources and have very low residual contamination.  Whereas, for the power spectrum analysis, we use the entire image map for the calculation, mixing both the good and the bad areas in the image cube.  Because of this difference, the power spectrum analysis requires a more stringent calibration accuracy of $\\sigma_a \\approx 0.05\\%$ compared to $\\sigma_a \\approx 0.2\\%$ for direct imaging.  This problem is most pronounced for the angular power spectrum analysis since the residual contamination is dominated by angular power.  It is mitigated partially in the 1D spherically-averaged power spectrum by the inclusion of spectral information--with its lower residual contamination power--in the calculation, and also in the 2D power spectrum. However, in the case of the 1D spherically-averaged power spectrum, it should be noted that the calibration accuracy is determined on the basis of a $\\sim 300$~hrs of integration by MWA-512. If we increase the amount of observing time, the requirement on the accuracy of calibration can be lowered because our model allows the calibration error to average toward zero with additional nights of observations.  The 2-D power spectrum, in particular, addresses our second key question, showing clear advantages for separating the residual contamination from the desired signal through distinct localization of the respective contributions in the the $\\kperp$ and $\\kpar$ plane. The results for both source position errors and residual calibration errors indicate that at $\\kperp \\lesssim 0.05$~Mpc$^{-1}$, we are able to probe most of the $\\kpar$ scales where the \\hi signal is dominant over the residual errors. In the 1D power spectrum we see dominant contribution from the residual errors around $\\k \\sim 0.1$~Mpc$^{-1}$, which can be probed in the 2D power spectra along the $\\kpar$ axes.  Finally, our third key question was whether we might expect to build a template library of residual contamination errors in the power spectrum domain in order to facilitate interpretation of the final power spectrum.  The results of this work indicate that it will indeed be possible.  Further, the significant similarities between the 2D power spectra for the source position error case and the gain calibration error case suggest that there may be common and easily identifiable properties of bright source residual contamination that are largely independent of the specific error causing the contamination.  This will be a valuable tool for upcoming experiments.  We have found that the IMLIN polynomial subtraction step is crucial not only for faint point source and diffuse foreground subtraction as studied in other works, but also for the success of the bright source foreground subtraction that we explored in this paper.  For the simulations included in this paper, we have performed the foreground subtraction of bright sources from a data-set of a minimum of 6~hours of observation (extrapolated to 300 and 5000 hours) in order to have the full effect of earth rotation synthesis.  However, the MWA may perform much of its bright source removal over much shorter timescales ($\\sim10$~minutes or less) as part of its real-time calibration pipeline.  The major implication for shorter time-scale removal of the foregrounds would be to break our assumption when modeling source position errors that the position errors are constant for the entire observation.  We also made the assumption in this paper that each antenna's calibration errors are perfectly correlated for an entire 6-hour observation night, but uncorrelated between observing nights.  If the residual calibration error where instead perfectly random between {\\it every} 8-second cycle of the real-time calibration, then we estimate it could be possible to achieve the desired residual contamination noise level and detect the redshifted 21~cm $\\hi$ signal from reionization with a significantly larger calibration error of $\\sigma_a~\\approx~2.5~\\%$. We have not performed our detailed simulation under this assumption, however, nor have we used the exact parameters that will be employed for the real-time calibration pipeline of the MWA.  The results from the angular power spectrum puts more stringent contraints on the accuracies in source position and calibration. Using a larger chunk ($>~ 125 kHz$)of frequency width might have reduced the constraints. However, we conclude that if a wider bandwidth is available, it is more advantageous to perform a 1D spherically averaged power spectrum than the angular power spectrum. This is due to the inclusion of the $\\kpar$ axes contribution to the power spectrum.  Lastly, we would like to emphasis that similar constraints can also be derived for other upcoming arrays, such as LOFAR and PAPER, as well as for future arrays like the Square Kilometer Array or a lunar array. But detailed simulations with the unique array specifications for each instrument would be required, which is beyond the scope of this paper.  We expect to build on our present analysis in future work by exploring other arrays, addressing the modified scenarios described above, and including additional calibration issues such as wide-field gain calibration of the primary beam and ionosphere."
    },
    "1005/1005.1148_arXiv.txt": {
        "abstract": "{\\it Spitzer} IRAC images (3.6, 4.5, 5.8 and 8.0$\\mu$m) and photometry of the star forming region M8 are presented. IRAC photometry reveals ongoing star formation in the M8 complex, with 64 Class 0/I and 168 Class II sources identified in several locations in the vicinity of sub-mm gas cores/clumps. Nearly 60\\% of these YSOs occur in about 7 small clusters. The spatial surface density of the clustered YSOs is determined to be about 10-20 YSOs/pc$^{2}$. Fresh star formation by the process of ``collect and collapse'' might have been triggered by the expanding HII regions and winds from massive stars. IRAC ratio images are generated and studied in order to identify possible diagnostic emission regions in M8. The image of 4.5$\\mu$m/8.0$\\mu$m reveals Br$\\alpha$ counterpart of the optical Hourglass HII region, while the ratio 8.0$\\mu$m/4.5$\\mu$m indicates PAH emission in a cavity-like structure to the east of the Hourglass. The ratio maps of  3.6$\\mu$m/4.5$\\mu$m, 5.8$\\mu$m/4.5$\\mu$m and 8.0$\\mu$m/4.5$\\mu$m seem to identify PAH emission regions in the sharp ridges and filamentary structures seen East to West and NE to SW in M8 complex. ",
        "introduction": "\\label{sec:intro} Messier 8 (M8), the Lagoon Nebula (or NGC 6523) is a well known galactic HII region (comprehensively reviewed in \\citet{Tothill08}, and the references therein) situated at a distance of 1.25 kpc \\citep{Arias06} in Sagittarius-Carina spiral arm of the Galaxy. The core of M8 contains a spectacular blister-type HII region, called the Hourglass nebula, ionized by the O7.5 V star Herschel 36 (Her 36) \\citep{Woolf61}. The HII region is embedded within a giant molecular cloud that extends eastwards to the young star cluster NGC 6530 of age 2 $\\times$ 10${^6}$ yrs \\citep{Lada76}. While Her 36 is responsible for the Hourglass and the ionised bubble surrounding it, the other early type stars in M8 complex, 9 Sgr (O4 V(f)) and HD165052(O6.5V+O7.5V) are believed to account for the ionised regions east of the central core/bubble \\citep{Goudis76,Lada76,Lynds82,Woodward86}. \\citet{Allen86} discovered a few near-IR sources in the vicinity of Her 36, designated by \\citet{Woodward90} as KS 1 to KS 5. From a high resolution near-IR study, \\citet{Arias06, Arias07} found the existence of a very young star cluster around Her 36, having an age of $\\sim$ 10${^6}$ yrs. \\citet{Barba07} discovered a number of HH objects in M8, which confirms by implication, the existence of very young stars undergoing the accretion phase of formation. Narrow-band imaging by the Hubble Space Telescope (HST) revealed the presence of proplyds in the neighbourhood of Her 36 \\citep{Stecklum98}. The core of M8 is detected by the Mid-course Space Experiment (MSX) in mid-infrared as a luminous extended source \\citep{Crowther03}. Color composite map of M8 in mid-infrared bands of {\\it Spitzer} Infrared Array Camera (IRAC) shows a ridge extending in east-west direction to the south-east of the Hourglass \\citep{Tothill08}. There appear also a number of filamentary structures extending in the NE-SW direction to the east of Her 36 or the Hourglass. Among the new compact star forming regions in the M8 complex, M8E stands out with its compact HII region powered by an early B type star (M8E-Radio) \\citep{Lada76,Wright77,Brand78,Mitchell91,Linz09}. \\citet{White97} discovered very intense CO line emission in mm and sub-mm wavelength regions from the central core of M8. Later, from a larger survey in mm and sub-mm continuum and CO lines, \\citet{Tothill02} found bright rims and dark lanes stretching in the east-west direction. \\citet{White97} found from CO (J=3--2) line observations a loose bipolar structure extending NW-SE from Her 36; while \\citet{Stecklum95} found the presence of a jet-like object very close to Her 36. These observations provide evidence for outflow activity around Her 36 region. As for the spectral diagnostics in the infrared region, \\citet{Woodward86} detected Br$\\gamma$ (2.17 $\\mu$m), Br$\\alpha$ (4.05 $\\mu$m) as well as Pf$\\epsilon$ (3.03 $\\mu$m) and the 3.28 $\\mu$m polycyclic aromatic hydrocarbons (PAH) emission to the west of Her 36. \\citet{Burton02} observed H$_{2}$ 1--0 S(1) line at 2.12 $\\mu$m near the Hourglass/Her 36 region attributed primarily to shock-excited molecular gas; but UV excitation can not be ruled out.  The M8 region seems to be quite complex and very interesting, owing to the presence of stellar winds and expanding HII region bubbles from massive stars, which can trigger fresh star formation by sweeping up and compressing the local dense interstellar matter. The afore-mentioned near-IR surveys studied the regions around Her 36 and NGC 6530, leaving the ridge regions far-east and south-east of Her 36 relatively under-explored in near-IR (except 2MASS survey) and mid-IR regions (except MSX survey). In view of this, we wish to examine in detail the extended region of M8 in the near-/mid-infrared images provided by {\\it Spitzer}, which are so far not looked into.  {\\it Spitzer} IRAC provides an opportunity with an unprecedented high spatial resolution in thermal infrared wavelength regime that is very useful in identifying embedded sources in massive star forming regions. IRAC has four wavelength bands (with $\\lambda_{eff}/\\Delta\\lambda$, 3.55/0.75, 4.49/1.0, 5.73/1.43 and 7.87/2.91 $\\mu$m) which include molecular emissions such as those from H$_{2}$ and PAH molecules, as well as recombination lines from hydrogen. The aims of the present study are to identify the embedded sources using the IRAC bands in order to classify the different stages of their evolution; to use ratio maps of the four IRAC bands in order to identify possible H$_{2}$, PAH or H emission regions, following the suggestions of \\citet{Smith05} and \\citet{Povich07}.  In Section 2, we describe the data used for the present study and the analysis tasks utilised. Section 3 presents the results and discussion on {\\it Spitzer} IRAC photometry of embedded sources associated with M8 complex. In this section, we also present the results and discussion on the ratio maps. In Section 4, we give the conclusions. ",
        "conclusions": "The important conclusions of this work are as follows:  \\begin{enumerate} \\item{{\\it Spitzer} IRAC photometry of M8 region revealed 64 Class 0/I and 168 Class II YSOs. About 60\\% of these are present in about 7 small clusters with spatial surface densities of 10-20 YSOs/pc$^{2}$;} \\item{These sources are positioned close to the sub-mm gas clumps and the filamentary or pillar like structures present in M8. It is possible that the formation of these sources could have been triggered by stellar winds or expanding HII regions associated with the massive stars in the region;} \\item{The ratio map Ch2/Ch4 reveals Br$\\alpha$ emission corresponding to the Hourglass HII region powered by Her 36 and its inverted ratio (Ch4/Ch2) identifies PAH emission in a cavity, east of the Hourglass;} \\item{The ratio maps Ch1/Ch2, Ch3/Ch2 and Ch4/Ch2 indicate the presence of PAH emission in both the ridges oriented along E-W and NE-SW directions.} \\end{enumerate}"
    },
    "1005/1005.2288_arXiv.txt": {
        "abstract": "The {\\it Approach unifying spin and charges}~\\cite{n92n93n95n01n07,pn06,% hn00hn02hn03hn06hn07hn08,gmdn07}, assuming that all the internal degrees of freedom---the spin, all the charges and the families---originate in $d> (1+ 3)$ in only two kinds of spins (the Dirac one and the only one existing beside the Dirac one and anticommuting with the Dirac one), is offering a new way in understanding the appearance of the families and the charges (in the case of charges the similarity with the Kaluza-Klein-like theories must be emphasized). A simple starting action in $d>(1+3)$ for gauge fields (the vielbeins and the two kinds of the spin connections) and a spinor (which carries only two kinds of spins and interacts with the corresponding gauge fields) manifests after particular breaks of the starting symmetry the massless {\\it four}  (rather than three) families with  the properties as  assumed by  the {\\it Standard model} for the three known families, and the additional four massive families. The lowest of these additional four families is stable. A part of the starting action contributes, together with the vielbeins, in the break of the electroweak symmetry manifesting in  $d=(1+3)$ the Yukawa couplings (determining the mixing matrices and the masses of the lower four families of fermions and influencing the properties of the higher four families) and the scalar field, which determines the masses of the gauge fields. The fourth family might be seen at the LHC, while the stable fifth family might be what is observed as the dark matter. ",
        "introduction": "\\label{introduction} The {\\it Standard model of the electroweak and colour interactions} (extended by the right handed neutrinos), fiting with around 25 assumptions and parameters all the existing experimental data, leaves  unanswered many open questions, among which are the questions about the origin of the charges ($U(1), SU(2), SU(3)$), of the families, and correspondingly of the Yukawa couplings of quarks and leptons and of the Higgs mechanism. Answering the question about the origin of families and their masses seems to me the most promising way leading beyond the today's knowledge about the elementary fermionic and bosonic fields.  The {\\it Approach unifying spins and charges}~\\cite{n92n93n95n01n07,pn06,hn00hn02hn03hn06hn07hn08,gmdn07} does have a chance to answer the above mentioned open questions. The question is, of course, whether and when the author~\\footnote{I started the project named the {\\it Approach unifying spins and charges} fifteen years ago, proving alone or together with collaborators step by step that such a theory has a real chance to answer the open questions of the {\\it Standard model}. The names of the collaborators and students can be found in the cited papers. } (and collaborators) will succeed to prove, that this theory really manifests in the low energy regime as the effective theory, postulated by the {\\it Standard model}.  Let me briefly present first  i.) the starting assumptions of the {\\it Approach} (in this section), ii.) a brief explanation of what does it offer (section~\\ref{predictions}), iii.) the so far made proofs and the obtained positive results (section~\\ref{LHCdarkmatter})  and at the end the open problems which are studying or are  waiting to be studied (section~\\ref{kkandotherproblems}).   The {\\it Approach} assumes~\\cite{n92n93n95n01n07,pn06,hn00hn02hn03hn06hn07hn08,gmdn07} a simple action in $d=(1+13)$-dimensional space \\begin{eqnarray} S &=& \\int \\; d^dx \\;{\\mathcal L}_{f} +  \\int \\; d^dx \\;  {\\mathcal L}_{g} \\label{action} \\end{eqnarray} with the Lagrange density  for a spinor, which carries  in $d=(1+13)$ {\\it two  kinds of the spin} and {\\it no charges}. One kind of spin is {\\it the Dirac spinor} in $(1+(d-1))$-dimensional space which, like in the Kaluza-Klein-like theories, takes care of the spin in $d=(1+3)$ and all the charges, and the {\\it second kind}, the new one,  which takes care of the families~\\footnote{This is the only theory in the literature to my knowledge, which does not explain the appearance of families by just postulating their numbers on one or another way, through the choice of a group, for example, but by offering the mechanism for generating families.}. There exist only two kinds of spins, corresponding to the left and the right multiplication of the Clifford algebra objects. The two kinds of spins are represented by the two kinds of the Clifford algebra objects~\\cite{n92n93n95n01n07,pn06,hn00hn02hn03hn06hn07hn08,gmdn07} \\begin{eqnarray} \\label{clifford} S^{ab}&=& \\frac{i}{4} (\\gamma^a \\gamma^b - \\gamma^b \\gamma^a),\\quad \\tilde{S}^{ab} = \\frac{i}{4} (\\tilde{\\gamma}^a \\tilde{\\gamma}^b - \\tilde{\\gamma}^b \\tilde{\\gamma}^a), \\nonumber\\\\ \\{\\gamma^a, \\gamma^b \\}_{+}&=& 2 \\eta^{ab}= \\{\\tilde{\\gamma}^a, \\tilde{\\gamma}^b \\}_{+}, \\quad \\{\\gamma^a, \\tilde{\\gamma}^b \\}_{+}=0,\\quad \\{S^{ab},\\tilde{S}^{cd}\\}_{-}= 0. \\end{eqnarray} The spinor interacts correspondingly only with the vielbeins $f^{\\alpha}{}_{a}$ and the two kinds of the spin connection fields, $ \\omega_{ab\\alpha}$ and  $\\tilde{\\omega}_{ab\\alpha}$, \\begin{eqnarray} {\\mathcal L}_f &=& \\frac{1}{2} (E\\bar{\\psi} \\, \\gamma^a p_{0a} \\psi) + h.c. \\nonumber\\\\ p_{0a }&=& f^{\\alpha}{}_a \\, (p_{\\alpha}  - \\frac{1}{2}  S^{ab}   \\omega_{ab\\alpha} - \\frac{1}{2}   \\tilde{S}^{ab}   \\tilde{\\omega}_{ab\\alpha}) + \\frac{1}{2E}\\, \\{p_{\\alpha}, f^{\\alpha}{}_{a} E \\}. \\label{lspinor} \\end{eqnarray} Correspondingly there is the Lagrange density for the gauge fields,  assumed to be linear in the curvature \\begin{eqnarray} \\label{lgauge} {\\mathcal L}_{g} &= &E  \\;(\\alpha \\, R + \\tilde{\\alpha}  \\tilde{R}), \\nonumber\\\\ R &=& f^{\\alpha [a}f^{\\beta b]} \\;(\\omega_{a b \\alpha,\\beta} - \\omega_{c a \\alpha} \\omega^{c}{}_{b \\beta}), \\tilde{R} = f^{\\alpha [a} f^{\\beta b]} \\;(\\tilde{\\omega}_{a b \\alpha,\\beta} - \\tilde{\\omega}_{c a \\alpha} \\tilde{\\omega}^{c}{}_{b \\beta}) \\end{eqnarray} and it is also the torsion field $ {\\cal T}^{\\beta}{}_{ab}= f^{\\alpha}{}_{[a} (f^{\\beta}{}_{b]})_{, \\alpha} + \\omega_{[a}{}^{c}{}_{b]} f^{\\beta}_{c}$. Indices $a$ and $\\alpha $ define the index in the tangent space and the Einstein index, respectively, $m$ and $\\mu $ determine the \"observed\" $(1+3)$-dimensional space, $s$ and $\\sigma$ the indices of \"non observed\" dimensions.    The action~Eq.(\\ref{action}), which manifests the families (through $\\tilde{S}^{ab}$) as the equivalent representations to the spinor representation of $S^{ab}$ (as can be seen in the last term of Eq.~(\\ref{clifford})), has a real chance to manifest after the spontaneous breaks of symmetries in the low energy regime,  as I shall comment bellow (sections~\\ref{predictions},~\\ref{LHCdarkmatter}), all the properties required by the {\\it Standard model} in order to fit the so far observed data. The {\\it Approach } does not only explain, for example, why  the left handed spinors carry the weak charge while the right ones do not, and predicts that for the spontaneous breaks of symmetries the vielbeins together with both kinds of the spin connection fields are responsible, as they  are also for the Yukawa couplings and the appearance of the scalar fields, but explains the mechanism for the appearance of families, predicting the fourth family to be possibly seen at the LHC (or at somewhat higher energies)~\\cite{pn06,gmdn07,nB09} and the stable fifth family, whose neutrinos and baryons with masses several hundred TeV/$c^2$ might explain the appearance of the dark matter~\\cite{gn09}.   The  {\\it Approach} confronted and still confronts several problems (among them are the problems common to all the Kaluza-Klein-like theories which  have vielbeins and spin connections as the only gauge fields, without  additional gauge fields in the bulk), which we are studying step by step when searching for possible ways of spontaneous breaking of the starting symmetries,  leading to the properties of the observed families of fermions and gauge and scalar fields, and looking for predictions the  {\\it Approach} is offering~\\cite{hn00hn02hn03hn06hn07hn08,nB09,gn09}.  I kindly ask the reader to look for more detailed presentation of the {\\it Approach} in the refs.~\\cite{pn06,gmdn07,nB09} and the references therein. ",
        "conclusions": "\\label{conclusion} I demonstrated in my talk that the {\\it Approach unifying spin and charges}, which assumes in $d=(1+13)-$dimensional space a  simple action for a gravitational field (manifesting with the spin connections and the vielbeins) and massless fermions carrying only {\\it two kinds of the spin} (the one presented by the Dirac matrices and the additional one anticommuting with the Dirac one--there is no the third kind of the spin), no charges, shows a new way beyond the {\\it Standard model of the electroweak and colour interactions}. It namely, for example, explains the origin of families, of the Yukawa couplings, of the charges, the appearing of the corresponding gauge fields and of the (two kinds of) scalar fields, it explains why only the left handed spinors carry the weak charge while the right handed ones do not, where does the dark matter originate, and others. The action  manifests at low (observable) energies after particular breaks of symmetries two times four families, with no Yukawa couplings (in comparison with the age of the universe) among the lower and the upper four families.    The {\\it Approach} predicts  the fourth family, to be possibly observed at the LHC or at somewhat higher energies and the stable fifth family, whose baryons and neutrinos  form the dark matter. I discussed briefly the properties of the lower four families, estimated at the tree level. Following the history of the fifth family members in the expanding universe up to today and estimating also the scattering properties of this fifth family on the ordinary matter, the evaluated masses of the fifth family quarks,  under the assumption that the lowest mass fifth family baryon is the fifth family neutron, are in the interval \\begin{eqnarray} \\label{all} 200 \\,{\\rm TeV} < m_{q_5}\\,c^2 <  10^5 \\,{\\rm TeV}. \\end{eqnarray} The fifth family neutrino mass is estimated to be in the interval: a few TeV $< m_{\\nu_5}\\,c^2 <$ a few hundreds TeV. \\subsection{Problems to be solved} \\label{kkandotherproblems} In the sections~(\\ref{conclusion},\\ref{introduction}) the promising sides of the {\\it Approach unifying spin and charges} are briefly over viewed and in the rest of talk the achievements presented. One of the conclusions we can make is that--since there are two kinds of the Clifford algebra objects and not only one (successfully used by Dirac 80 years ago) describing the spin of fermions and in the {\\it Approach}  the spin and the charges, and since the other generates the equivalent representations with respect to the first one, while the families are the equivalent representations with respect to the spin and the charges--the second kind of the spin can be used to describe families, or even must, since if not, we should explain, why only one kind of the spin manifests at the low energy regime.  There are, however, several questions which should be solved  before accepting the {\\it Approach} as the right way beyond the {\\it Standard model}, offering the  right explanation  for the origin of families, charges and gauge fields. Since the {\\it Approach} assumes that the dimension of the space time is larger than (1+3), like do the Kaluza-Klein-like theories as well as the theories with strings and membranes, it must be asked what is at all the dimension of the space-time. What does cause or trigger the spontaneous breaks of symmetries in the evolution of the universe? What does determine the phase transitions?  We put a lot of efforts in understanding the properties of higher dimensional spaces and in understanding the properties of fermions after breaking symmetries. The reader can find more about our understanding and the proposed solutions in the references~\\cite{hn00hn02hn03hn06hn07hn08}. It is  hard to evaluate how do the spontaneous (non adiabatic) breaks occur, what causes them and how do they influence the properties of all kinds of fields. To understand better the differences in the properties of a family members, we are estimating their properties bellow the tree level. Is it the expectation that there are effects  bellow the tree level which are responsible for the differences in properties of family members (although some differences manifest already on the tree level) correct? Will such calculations show that the fifth family $u$-quark is heavier than the $d$-quark  mainly due to the repulsive electrostatic contribution ($\\approx 100$ GeV)?  The estimation of the behaviour of the coloured fifth family clusters during the colour phase transition  leads to the conclusion that the fifth family quarks either annihilate or form the fifth family neutrons or antineutrons. Will more accurate evaluations confirm this estimation? How many of the fifth family neutrinos and neutrons annihilate in the  interval bellow the weak phase transition due to possibly large weak cross section for the annihilation in this region?   Many a question presented here are under consideration, many of them not even written here wait to be studied. I invite the audience to contribute to next steps of (to my understanding) a really promising way beyond the {\\it Standard model}."
    },
    "1005/1005.0843_arXiv.txt": {
        "abstract": "We perform a detailed study of the sensitivity to the anisotropies related to Dark Matter (DM) annihilation in the Isotropic Gamma-Ray Background (IGRB) as measured by the Fermi Large Area Telescope (Fermi-LAT). For the first time, we take into account the effects of the Galactic foregrounds and use a realistic representation of the Fermi-LAT. We implement an analysis pipeline which simulates Fermi-LAT data sets starting from model maps of the Galactic foregrounds, the Fermi resolved point sources, the extra-galactic diffuse emission and the signal from DM annihilation. The effects of the detector are taken into account by convolving the model maps with the Fermi-LAT instrumental response. We then use the angular power spectrum to characterize the anisotropy properties of the simulated data and to study the sensitivity to DM. We consider DM anisotropies of extra-galactic origin and of Galactic origin (which can be generated through annihilation in the Milky Way sub-structures) as opposed to a background of anisotropies generated by sources of astrophysical origin, blazars for example. We find that with statistics from 5 years of observation Fermi is sensitive to a DM contribution at the level of 1\\%-10\\% of the measured IGRB depending on the DM mass $m_\\chi$ and annihilation mode. In terms of the thermally averaged cross section $\\left< \\sigma_{\\rm{A}}v\\right>$, this corresponds to $\\sim 10^{-25}$cm$^{3}$s$^{-1}$, i.e. slightly above  the typical expectations for a thermal relic, for low values of the DM mass $m_\\chi\\alt 100$ GeV. The anisotropy method for DM searches has a sensitivity comparable to the usual methods based only on the energy spectrum and thus constitutes an independent and complementary piece of information in the DM puzzle. ",
        "introduction": "Combined studies of the  cosmic microwave background  radiation, supernova cosmology and large Galaxy redshift surveys provide nowadays compelling evidence of the existence of non-baryonic Dark Matter (DM) \\cite{Komatsu:2008hk,Tegmark:2006az}. Consistency with the observed structure of the Universe, especially at large (above galactic) scales favors Cold Dark Matter (CDM), i.e. the particles constituting the cosmologically required dark matter have to be moving non-relativistically. This property is always fulfilled by particles of mass in the GeV to TeV region that interact with the weak interaction strength, i.e. weakly interacting massive particles (WIMPs). WIMPs can annihilate or decay to detectable standard model particles, in particular gamma-rays, giving rise to the possibility of ``indirect detection'' \\cite{Bertone:2004pz,Jungman:1995df,Bergstrom:2000pn,Baltz:2008wd}. A  standard approach is to look for spectral signatures in the energy spectrum of gamma-rays from celestial objects, either targeting the continuum spectrum produced by decay of secondary pions or the more feeble signature of a line due to DM annihilating directly into photons.  Another possible experimental signature is given by differences in the spatial distribution of gamma-ray signals induced by DM as compared to conventional astrophysical sources. In particular, the small scale fluctuations in an otherwise isotropic gamma-ray background (IGRB) might be different in the case of annihilating DM, since its flux scales as the squared density. After the pioneering study of Ando and Komatsu~\\cite{Ando:2005xg} , followed by two other more detailed works \\cite{Ando:2006mt,Ando:2006cr}, the issue of DM annihilation and anisotropies in the isotropic gamma background has raised increasing interest. In connection to the extra-galactic gamma-ray background,  anisotropy has  been further studied in \\cite{Cuoco:2006tr,Cuoco:2007sh,Cuoco:2008zz} while predictions from the Millenium II simulation have recently been described in Zavala et al. (2009). It has also been realized that Galactic substructures can produce an almost isotropic gamma background  with similar fluctuation properties as the extragalactic background and with promising chances of detection \\cite{SiegalGaskins:2008ge,Ando:2009fp,SiegalGaskins:2009ux,Hensley:2009gh}; see also \\cite{Pieri:2007ir,Pieri:2009je}). Other works have investigated the anisotropy pattern resulting from both the Galactic and extragalactic contribution \\cite{Fornasa:2009qh,Hooper:2007be,Ibarra:2009nw} or from the the DM annihilation around intermediate mass black holes \\cite{Taoso:2008qz}. Anisotropies from DM annihilation appearing in the radio sky have been  investigated in \\cite{Zhang:2008rs}. Finally,-- apart from the power spectrum -- an interesting approach is to investigate the probability distribution of the fluctuations in order to detect possible departures from Gaussianity (or, rather, from Poisson statistics, in the common situation of low number of counts), thereby providing another handle on the separation of conventional astrophysical and DM contribution \\cite{Lee:2008fm,Dodelson:2009ih}. Anisotropy properties of astrophysical processes other than DM contributing to the isotropic gamma background include resolved point sources \\cite{Ando:2006mt}, galaxy clusters and shocks \\cite{Miniati:2007ke,Keshet:2002sw} and emission from normal galaxies \\cite{Ando:2009nk}.   The Fermi Large Area Telescope (Fermi-LAT), which was launched June 11, 2008, is the most suitable tool for these kind of studies due to its large field of view ($\\sim$ 2.4 sr), energy coverage from 100 MeV to 300 GeV \\cite{Atwood:2009ez} and very good angular resolution. We will thus focus in the following on the prospect to detect DM anisotropies with Fermi-LAT.   All of the above studies assume either that the impact of Galactic foregrounds on the DM induced anisotropies is negligible or that the Galactic foregrounds can be cleaned at a suitable level to allow DM anisotropies studies.  Further, the effects of the Fermi-LAT detector are generally taken into account only in an approximate way, though a detailed representation of the Fermi instrument is publicly available.  The aim of this work is to address the above issues in detail. In particular our approach will be the following: we include the Galactic foregrounds explicitly using a model based on the GALPROP code \\cite{Strong:1998pw,Moskalenko:1997gh,Strong:2004de} and we then use this model to build masks to select the most promising regions in the sky suitable for the analysis of the anisotropies. In these regions we study the anisotropies in the scenario in which a DM contribution is present over the Galactic foregrounds and the extra-galactic emission from blazars compared to the case without DM and we derive the sensitivity to the DM signal employing anisotropy studies with Fermi-LAT data. Both the cases of Galactic DM in substructures and extra-galactic DM are considered separately. The effect of the Fermi-LAT response, including the effects of the point spread function, the effective area, as well as non-uniformities in exposure, are included in all calculations at a realistic level by using the instrument response functions provided by the Fermi Science Tools.  This approach is suitable for a sensitivity study like the one presented here, and does not aim to fully reproduce a real data analysis for which foreground cleaning (not performed in this study) most likely will be required. One of the findings of our analysis is indeed that even with a very generous masking, foregrounds still have a sizeable contribution to the overall anisotropies and represent an important limitation to DM sensitivity. This implies that foreground cleaning will be a crucial step in the data analysis pipeline. On the other hand, if cleaning can be achieved in an effective way, this also means that the sensitivity to DM can be further improved with respect to the results derived in this paper. Our results can thus be regarded as conservative from this point of view. Foreground cleaning will be the subject of forthcoming investigations.   The paper is structured as follows. In section II we introduce and discuss the various components contributing to the Galactic and extragalactic gamma-ray emission. In sections III and IV we review their anisotropy properties and we describe the tools and methods for their analysis. In section V we  briefly describe the Fermi-LAT instrument and the characteristics relevant for our analysis. In section VI we discuss the results of our data simulation pipeline. Finally, in section VII we evaluate Fermi's sensitivity to distinguish a scenario in which the diffuse gamma-ray emission has a DM contribution (either extragalactic or galactic in origin) with respect to the case in which DM is absent. Conclusions are presented in section VIII.     \\begin{figure*} \\vspace{-0.0pc} \\begin{center} \\vspace{-0.5pc}\\includegraphics[width=1.90\\columnwidth,angle=0]{DiffPlot.eps} \\end{center} \\vspace{-1.0pc} \\caption{Energy spectrum of the various components contributing to the gamma sky above $30^\\circ$ (left) and above $60^\\circ$ (right) galactic latitude. The isotropic component is the determination from Abdo et al. (2009b). Also shown is an example of extragalactic WIMP emission normalized with the model of Ullio et al. (2002) (with NFW halo profile and no halo substructures). The spectrum of the resolved point sources is calculated from the Fermi 11-month source catalogue (Ballet 2009). \\label{fig:diffuse}} \\end{figure*} ",
        "conclusions": "In the present analysis we have calculated the sensitivity to Dark Matter annihilation exploiting the characteristic anisotropy properties that DM would imprint in the Extra-Galactic Gamma Background. We performed a realistic simulation of the Fermi-LAT including the effect of the Point Spread Function, effective area and of the anisotropic sky exposure.  More importantly, we investigated the impact of Galactic foregrounds on the DM anisotropy sensitivity using a GALPROP model to produce the sky maps of the Galactic gamma emission. The contribution of Galactic foregrounds is generally non negligible even above Galactic latitudes of about $60^\\circ$. Foreground cleaning will thus be an important (and likely challenging) issue in the analysis of the real data. For the present work we do not attempt foreground cleaning (except for masking the Galactic plane and point sources) and we only require that the anisotropy from DM has to exceed the level of foregrounds anisotropy to be detectable. Further, we only analyze a high Galactic latitude region where the model dependence of the foreground is expected to be moderate. In this respect our analysis can be regarded as a conservative one: In reality, if foreground cleaning can be effectively implemented, a sensitivity to a lower level of anisotropy can be achievable. Also, a larger fraction of the sky can be employed than the small conservative one used in the present analysis. We  leave a more careful investigation of the issue for future work.  Despite the difficulties involved in the foreground analysis, we find  that the prospects for detecting DM anisotropy signatures are promising. For some WIMP masses, sensitivity to  $\\left< \\sigma_{\\rm{A}}v\\right> \\sim 3\\times10^{-26}$cm$^{3}$s$^{-1}$(i.e. the thermal relic value) can be achieved using the anisotropy information alone. Interestingly, these constraints are comparable to what is achievable using the energy spectrum alone \\cite{Abdo:2010dk} . The anisotropy information is thus complementary to the energy spectrum and provides a cross-check and an increase in the overall sensitivity to the DM signal.  We also investigated the recently proposed scenario of a population of DM substructures in the Milky Way halo as predicted by N-body simulations. These subhaloes give rise to an  annihilation signal which is almost uniform in the sky and which can be comparable in intensity to the observed isotropic gamma background. For this case,  the anisotropy signal is predicted to be even higher than for the extragalactic case due to the huge clumpiness of the DM distribution that correspondingly gives an enhanced sensitivity. We find, indeed, that when increasing the anisotropy the sensitivity to the DM signal increases as well with an improvement of a factor of $\\sim$10 for $m_\\chi\\agt100$ GeV.  In summary, simulations indicate that ,with the sensitivity of the Fermi-LAT, detailed anisotropies studies of the gamma-ray sky can be performed. The next obvious step will be the analysis of the actual data. Anisotropy analyses will soon provide a complementary source of information which will give further insights into the DM problem and the properties of high-energy astrophysical sources."
    },
    "1005/1005.5711_arXiv.txt": {
        "abstract": "We study the effect of dark matter (DM) particles in the Sun, focusing in particular on the possible reduction of the solar neutrinos flux due to the energy carried away by DM particles from the innermost regions of the Sun,  and to the consequent reduction of the temperature of the solar core. We find that in the very low-mass range between 4 and 10 GeV, recently advocated to explain the findings of the DAMA and CoGent experiments, the effects on neutrino fluxes are detectable only for DM models with very small, or vanishing, self-annihilation cross section, such as the so-called asymmetric DM models, and we study the combination of DM masses and Spin Dependent cross sections which can be excluded with current solar neutrino data. Finally, we revisit the recent claim that DM models with large self-interacting cross sections can lead to a modification of the position of the convective zone, alleviating or solving the solar composition problem. We show that when the `geometric' upper limit on the capture rate is correctly taken into account, the effects of DM are reduced by orders of magnitude, and the position of the convective zone remains unchanged. ",
        "introduction": "The identification of Dark Matter (DM) is one of the most important open problems of modern cosmology and particle physics. Among the many DM candidates proposed in the literature, Weakly Interacting Massive Particles (WIMPs) are the most popular, as they naturally achieve the appropriate relic density and naturally arise in extensions of the Standard Model of particle physics that will soon be probed in accelerators. WIMPs are currently searched for in a variety of experiments (see Ref.~\\cite{reviews} for recent reviews on particle DM, including a discussion of ongoing direct, indirect and accelerator searches).  Here, we focus on the effect of WIMPs on stars, and more specifically on the Sun. In general, WIMPs crossing a star can in principle scatter off baryons to velocities lower than the star escape velocity and be therefore gravitationally trapped by the celestial body, the efficiency of the process depending on the WIMPs scattering cross section and WIMPs ambient density. Once captured, WIMPs can still scatter with the star nuclei and therefore transfer energy inside the object. In addition, WIMPs can eventually annihilate, providing therefore an exotic source of energy. These effects have been investigated in the past, in particular in the context of WIMPs 'cosmions' in order to solve the solar neutrino problem \\cite{cosmions,Dearborn:1990mm}. This solution has then been discarded in favour of neutrino oscillations and the combination of the cosmion scattering cross-sections and masses suggested by the solar neutrino problem is nowadays excluded by direct detection constraints.  Still, the intriguing possibility that WIMPs modify the properties of the stars has recently prompted new interest on the subject. Recent works have focused in particular on compact objects ~\\cite{Moska, Bertone:2007ae}, main sequence stars at the Galactic center ~\\cite{Fairbairn08,Scott:2008ns} and first stars \\cite{Ascasibar07,Freese,Freeseb,Iocco2,Freese2,Freese3,Iocco:2008rb,Taoso:2008kw,Yoon:2008km, RipamontiDM-star}. In all these studies the celestial objects are placed in environments with high DM density, thus enhancing the numbers of WIMPs captured by the star and therefore the WIMPs transport and annihilation energy budgets. However, in these scenarios it is not possible to set robust constraints on the DM cross sections and mass because of the large uncertainties on the DM densities and/or the lack of precise enough observations of the star targets.  The situation is quite different for the Sun since its properties have been measured with good precision. Because of that, the Sun can be used as a diagnostic tool to test small modifications of its structure induced by DM. The most recent works in this direction have focused on modifications of the solar neutrino fluxes and helioseismology data \\cite{Lopes:2001ig,Lopes:2002gp,Bottino:2002pd}. Here we perform a complete and self-consistent calculation of the Sun evolution inside the galactic DM halo. In particular we extend previous works considering light DM candidates, with masses in the range suggested by DAMA \\cite{Dama} CoGent \\cite{Aalseth:2010vx} and CDMSII \\cite{Ahmed:2009zw} experiments. We consider both standard annihilating WIMPs models and scenarios with negligible DM annihilations. We then move on Self Interacting Dark Matter (SIDM) models \\cite{Spergel:1999mh}, which has recently been invoked as a solution of the solar composition problem \\cite{Frandsen:2010yj}. Finally, we focus on Inelastic DM models \\cite{TuckerSmith:2001hy}.  The paper is organised as follows: in Section \\ref{sec:theoretical} we introduce the DM models used in our calculations. In Section \\ref{sec:code} we present the GENEVA stellar evolution code. In Sec \\ref{sec:diagnostictools} we discuss how the solar neutrino fluxes can be used to constrain the DM parameter space. We present our results for standard annihilating WIMPs in Section \\ref{sec:pheno}, for asymmetric DM models in Section \\ref{sec:results}, for SIDM models in Section \\ref{sec:self-inter} and for Inelastic DM in Section \\ref{sec:inelastic}. Finally we summarize our conclusions in Section \\ref{sec:concl}. Technical details and the formulas implemented in the code are in Appendix \\ref{sec:formalism}. ",
        "conclusions": "\\label{sec:concl}  By the use of a stellar evolution code, we have studied the modifications on the Sun structure induced by DM particles captured by the Sun during its lifetime. For standard WIMPs an equilibrium between the capture and the annihilation rates is reached in short timescales. After this transient period, the number of WIMPs inside the Sun is constant and we find it to be too small to produce any observable effect on the Sun properties.  However, for scenarios with null or very small annihilation cross sections, $\\sigv \\lesssim 10^{-33}$  cm$^3$s$^{-1},$ the number of trapped WIMPs is significantly increased. In this case, the transport of energy by the WIMPs from the interior of the Sun to the outer shells can dramatically reduce the core temperature, the most important consequence being a reduction of the $^8B$ flux. Considering the present theoretical and experimental uncertainties on $\\nuflub8$, we have studied the combination of DM masses and SD scattering cross sections which can be ruled our with this argument, finding only a small region of the parameter space which is not already excluded by direct detection bounds. Even with a significant decrease of the uncertainties on $\\nuflub8$, the region of the parameter space which can be probed remains approximately the same, therefore future experimental advances will not significantly change the situation.  We have also considered SIDM candidates which have been recently invoked as a solution of the solar composition problem. Correctly implementing the geometrical limit, we find that the number of WIMPs captured on the star is not sufficient to produce significant deviations on the Sun properties, unless one considers very high scattering cross sections, where however the models become in tension with direct detection bounds. Even in this case, we encounter significant modifications of the density and temperature profiles only in the inner regions of the Sun while the outer shells are not significantly affected. We can therefore exclude that the transport of energy produced by a population of WIMPs can solve or even alleviate the solar composition problem.  Finally, we notice that the WIMPs energy transport may produce dramatic effects on the structure of stars settled in high DM density environments, like the galactic center. We leave the investigation of this scenario for future work.  When this manuscript was ready for submission, a similar work appeared on the arXiv, focusing on the effects of WIMPs in the Sun \\cite{Cumberbatch:2010hh}. While our work mainly focus on the modifications of the solar neutrino fluxes, Ref. \\cite{Cumberbatch:2010hh} considers future helioseismology data as a diagnostic tool to constrain the WIMPs parameter space."
    },
    "1005/1005.3137_arXiv.txt": {
        "abstract": "{About two dozen exo-solar debris systems have been spatially resolved.  These debris discs commonly display a variety of structural features such as clumps, rings, belts, excentric distributions and spiral patterns. In most cases, these features are believed to be formed, shaped and maintained by the dynamical influence of planets orbiting the host stars. In very few cases has the presence of the dynamically important planet(s) been inferred from direct observation.} {The solar-type star \\qeri\\ is known to be surrounded by debris, extended on scales of \\lapprox\\,30\\asec. The star is also known to host at least one planet, albeit on an orbit far too small to make it responsible for structures at distances of tens to hundreds of AU. The aim of the present investigation is twofold: to determine the optical and material properties of the debris and to infer the spatial distribution of the dust, which may hint at the presence of additional planets.} {The photodetector array camera and spectrometer (PACS) aboard the {{\\it Herschel}} Space Observatory allows imaging observations in the far infrared at unprecedented resolution, i.e. at better than 6\\asec\\ to 12\\asec\\ over the wavelength range of 60\\,\\um\\ to 210\\,\\um. Together with the results from ground-based observations, these spatially resolved data can be modelled to determine the nature of the debris and its evolution more reliably than what would be possible from unresolved data alone.} {For the first time has the \\qeri\\ disc been resolved at far infrared wavelengths. The PACS observations at 70\\,\\um, 100\\,\\um\\ and 160\\,\\um\\ reveal an oval image showing a disc-like structure in all bands, the size of which increases with wavelength. Assuming a circular shape yields the inclination of its equatorial plane with respect to that of the sky, $i > 53$\\adeg. The results of image de-convolution indicate that $i$ likely is larger than 63\\adeg, where 90\\adeg\\ corresponds to an edge-on disc.} {The observed emission is thermal and optically thin. The resolved data are consistent with debris at temperatures below 30\\,K at radii larger than 120\\,AU. From image de-convolution, we find that \\qeri\\ is surrounded by an about 40\\,AU wide ring at the radial distance of \\about\\,85\\,AU. This is the first real Edgeworth-Kuiper Belt analogue ever observed.} ",
        "introduction": "We report new observations of the debris disc around \\qeri\\ (HD\\,10647, HR\\,506) using the {\\it Herschel} Space Observatory \\citep{pilbratt2010}. The observations form part of a larger Key Programme (KP), viz. DUNES\\footnote{DUst around NEarby Stars, {\\tiny {\\ttfamily http://www.mpia-hd.mpg.de/DUNES/}}}, which is described in more detail by \\citet{eiroa2010}. Here, we give a brief summary to put the contents of this Letter into context. The DUNES KP is a sensitivity limited study with the goal of discovering and characterising extra-solar analogues of the Edgeworth-Kuiper Belt (EKB) in an unbiased, statistical sample of nearby F, G and K main-sequence stars. The sample is volume limited, with distances \\lapprox\\,20\\,pc, and spans a broad range of stellar ages, from \\about\\,0.1 to roughly 10\\,Gyr. In addition to the object of the present study (\\qeri), a number of M- and A-type stars will be observed in collaboration with the DEBRIS-KP team \\citep{matthews2010}, implying that the whole sample covers a decade in stellar mass from 0.2 to 2\\,\\msun.  The PACS \\citep{poglitsch2010} observations at 100\\,\\um\\ aim at the detection of the stellar photospheres down to the confusion noise with a signal to noise ratio (S/N) of at least 5. Together with observations in the other {\\it Herschel} bands, this will lead to an unprecedented characterisation of discs and will allow detailed theoretical modelling. It is foreseen that fractional luminosities $L_{\\rm dust} / L_{\\odot}$ of a few times $10^{-7}$ will be reached, i.e. similar to that of the EKB of the Solar System \\citep{stern1996,jewitt2009} and a more than order of magnitude improvement over {\\it Spitzer} data \\citep{bryden2009}.  \\begin{figure*}[t] \\resizebox{\\hsize}{!}{ \\rotatebox{00}{\\includegraphics{14601fig1.eps}} } \\caption{PACS photometric imaging of \\qeri\\ at, from left to right, 70\\,\\um\\  ({\\it blue}), 100\\,\\um\\ ({\\it green}) and 160\\,\\um\\ ({\\it red}). The 70\\,\\um\\ image was taken in chop-nod mode, whereas the other two in scan-map mode. The upper panels display the reduced observations. Below, de-convolved images are shown, using observations of $\\alpha$\\,Boo for the definition of the PSF. Displayed are the results for ten iterations of a MEM algorithm \\citep{hollis1992}. The star defines the origin of the frames, i.e. offset coordinates (0, 0). Within the positional accuracy (2\\asec\\ rms), the stellar position and the centre of the elliptical brightness distributions coincide (see Table\\,\\ref{results}) and offsets are in seconds of arc. The lowest contours are at 5\\% of the maximum values and consecutive steps are also by this amount. At the distance of the star, 20\\asec\\ corresponds to 350\\,AU. } \\label{observed} \\end{figure*}  A main-sequence star with an infrared excess larger by more than three orders of magnitude than this limit is the late F-type star \\qeri. This will potentially allow to study the material giving rise to the excess in great detail. The star is known to be accompanied by a giant planet, \\qeri\\,b, orbiting at the distance of 2\\,AU \\citep[][and references therein] {butler2006}, which corresponds to \\asecdot{0}{1} at the distance of the system \\citep[17.35\\,pc;][]{perryman1997,vanleeuwen2007}. Such a small angle will not be resolved by the observations described in this paper, but \\qeri\\ is also surrounded by a ring or belt system of scattering and thermally emitting dust particles. The extent of the system, which from optical to sub-millimetre wavelengths is up to several tens of arc seconds in size \\citep[e.g.,][and references therein]{stapelfeldt2007,liseau2008}, should be readily accessible to PACS at the {\\it Herschel} telescope ($\\theta_{\\rm diff}^{\\,\\prime \\prime} = 7 \\times \\lambda_{\\mu{\\rm m}}/100\\,\\mu{\\rm m}$).  This contribution presents the observed properties of the \\qeri\\ system by {\\it Herschel}. The results of theoretical model calculations will be communicated by Augereau et al. (in prep.). The observations and data reduction are described in Sect.\\,2, with the results presented in Sect.\\,3. These are discussed in Sect.\\,4 and, finally in Sect.\\,5, our main conclusions are briefly summarised.  \\begin{table*} \\begin{flushleft} \\caption{\\label{results} Brightness, orientation and extent of the \\qeri\\ debris system} \\resizebox{\\hsize}{!}{ \\begin{tabular}{lcllllll} \\hline \\noalign{\\smallskip} Observation & $\\Delta {\\rm RA}$, $\\Delta {\\rm Dec}^{a}$ \t& Flux$^{b}$\t\t\t& 2a$^{b}$\t&   2b$^{b}$\t\t\t& PA$^{c}$     \t& $i^{\\,d}$\t\t& Alternate flux\t\t\\\\ ID and Mode    \t\t&\t(\\asec, \\asec)\t\t\t& (mJy)\t\t\t\t&(\\asec)\t\t&  (\\asec)\t\t\t\t& (\\adeg) \t\t& (\\adeg)\t\t& measurement (mJy)\t\\\\ \\noalign{\\smallskip} \\hline \\\\ 1342187142 ~~ - \\phantom{1}70 :  chop-nod&$+0.9$, $-0.9$& $ 828 \\pm 83$ \t\t& 13.9 (11.0)\t& \\phantom{1}7.9 (4.4)\t&   54.4 (60.0)\t& 55.4 (66.4)\t& \\phantom{1}$859 \\pm  \\phantom{1}6$ {\\it Spitzer}/MIPS$^{e}$\\\\ 1342187141~~~ - 100 :  chop-nod\t &   $+2$, $0$\t\t& $ 816 \\pm 82$\t\t& 15.5 (12.6)\t& \\phantom{1}8.8 (5.1) \t&   54.4 (56.2)\t& 55.4 (66.1)\t& $1080\\pm 36$ IRAS FSC\\\\ 134218739/40 - 100 : scan map\t&     $0$, $-1.4$\t& $ 810 \\pm 81$\t\t& 15.1 (12.5)\t& \\phantom{1}8.8 (4.7)\t&   56.1 (54.8)\t& 54.4 (67.9)\t& $1080\\pm 36$ IRAS FSC\\\\ 1342187142 ~~~- 160 :  chop-nod \t& $-1.4$, $+0.9$\t& $ 529 \\pm 106$\t\t& 18.9 (13.5)\t& 12.3 (7.8)\t\t\t&   51.3 (48.9)\t& 49.4 (54.7)\t& \\phantom{1}$453 \\pm 50$ {\\it Spitzer}/MIPS$^{f}$\\\\ 134218739/40 - 160 : scan map\t&$+1.4$, $-1.6$\t& $ 537 \\pm 107$\t\t& 19.3 (13.3)\t& 12.5 (6.5)\t \t\t&   51.2 (55.7)\t& 49.6 (60.7)\t& \\phantom{1}$453 \\pm 50$ {\\it Spitzer}/MIPS$^{f}$\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} } \\end{flushleft} Notes to the Table: \\\\ $^{a}$ Offsets of ellipse centre relative to stellar coordinates: proper motion corrected to IRCS 2009.89 \\citep[+0.186\\,s and \\asecdot{$-1$}{055}][]{perryman1997,vanleeuwen2007} \\radot{01}{42}{29}{502}, \\decdot{$-53$}{44}{28}{06}. The photospheric fluxes are 17\\,mJy  at 70\\,\\um, 8\\,mJy at 100\\,\\um\\ and 3\\,mJy at 160\\,\\um\\ (cf. Fig.\\,\\ref{SED}). \\\\ $^{b}$ Flux for Gaussian ellipse with fitted FWHM to major and minor axis, 2a and 2b, respectively. Values in parentheses refer to de-convolved images.\\\\ $^{c}$ Position angle measured from North over East.\\\\ $^{d}$ Lower limit to the inclination, where $i=90$\\adeg\\ refers to an edge-on geometry.\\\\ $^{e}$ \\citet{trilling2008}.\\\\ $^{ f}$ \\citet{tanner2009}. \\end{table*} ",
        "conclusions": "Based on imaging observations with PACS in the three photometric bands at 70\\,\\um, 100\\,\\um\\ and 160\\,\\um\\ we find that \\begin{itemize} \\item[$\\bullet$] the debris around the solar-type star \\qeri\\ has an oval-shaped brightness distribution, the size of which increases with the wavelength; \\item[$\\bullet$] the integrated flux density at these wavelengths leads to an SED which is in good agreement with earlier results; \\item[$\\bullet$] the very high signal-to-noise of the 100\\,\\um\\ scan map is adequate to sharpen the image using an image de-convolution technique, revealing a ring-like structure with maximum surface density at \\about\\,85\\,AU from the star; \\item[$\\bullet$] with a width of about 35 to 45\\,AU, this ring or belt around the F9\\,V star \\qeri\\ is similar to the Edgeworth-Kuiper Belt around the Sun. This may hint at the presence of another planet, \\qeri\\,c. \\end{itemize}  \\acknowledgement{We wish to acknowledge the support by the SNSB, the CNES, the PNP and the MICINN. In addition, we have benefited from HCSS / HSpot / HIPE which are joint developments by the {\\it Herschel} Science Ground Segment Consortium, consisting of ESA, the NASA {\\it Herschel} Science Center, and the HIFI, PACS and SPIRE consortia.}"
    },
    "1005/1005.1074_arXiv.txt": {
        "abstract": "{ The star formation rate (SFR) is a key parameter in the study of galaxy evolution. {The accuracy of SFR measurements at z$\\sim$2 has been questioned following a disagreement between observations and theoretical models.} The latter predict SFRs at this redshift that are typically a factor 4 or more lower than the measurements. We present star-formation rates based on calorimetric measurements of the far-infrared (FIR) luminosities for massive 1.5$<$z$<$2.5, normal star-forming galaxies (SFGs), which do not depend on extinction corrections and/or extrapolations of spectral energy distributions. The measurements are based on observations in GOODS-N with the Photodetector Array Camera \\& Spectrometer (PACS) onboard {\\it Herschel}, as part of the PACS Evolutionary Probe (PEP) project, that resolve for the first time individual SFGs at these redshifts at FIR wavelengths. We compare FIR-based SFRs to the more commonly used 24~$\\mu$m and UV SFRs. We find that SFRs from 24~$\\mu$m alone are higher by a factor of $\\sim$4--7.5 than the true SFRs. This overestimation depends on luminosity: gradually increasing for $log$~L(24~$\\mu$m)$>12.2$~L$_\\odot$. The SFGs and AGNs tend to exhibit the same 24$\\mu$m excess. The UV SFRs are in closer agreement with the FIR-based SFRs. Using a Calzetti UV extinction correction results in a  mean excess of up to 0.3~dex and a scatter of 0.35~dex from the FIR SFRs. The previous UV SFRs are thus confirmed and the mean excess, while narrowing the gap, is insufficient to explain the discrepancy between the observed SFRs and simulation predictions. } ",
        "introduction": "The evolution of massive galaxies around z$\\sim$2, when the cosmic star-formation rate density was at its peak, is currently the subject of intense study. A key parameter in constraining their nature and stellar-mass buildup is the instantaneous star-formation rate (SFR). While there are various methods for estimating SFRs \\citep[see ][ for a review]{Kennicutt98}, it is notoriously difficult to constrain SFRs of z$\\sim$2 galaxies {with an accuracy better than a factor of a few}. One of the reasons is that calibrations between observed luminosities (rest-frame UV, mid-IR, and submm being the most accessible and widely used) and bolometric SFRs are based on local galaxies and their validity at high z has not yet been firmly established.  Both UV (BX/BM) and optical/near-IR (BzK) selected samples of star-forming galaxies at z$\\sim$2 have been successful in producing large samples for galaxy evolution studies \\citep{Reddy05, Perez-Gonzalez05, Daddi07a, Santini09}. However, UV-based SFRs are very sensitive to dust obscuration and there are uncertainties about the applicable reddening laws and dust/source geometry at high~z \\citep[e.g., ][]{Reddy10}. Estimates from the 24~$\\mu$m luminosities rely on template-based extrapolations to the total IR luminosity that are poorly constrained at high~z and may be affected by an AGN. Likewise, submm luminosities rely on assumptions about IR SEDs, which depend on poorly known dust temperatures, composition, and spatial distribution.  Cosmological simulations predict a tight correlation between stellar mass and SFR with a slope similar to that observed \\citep[e.g.,][]{Finlator06}, but with SFRs that are a factor $\\sim$4 lower than those inferred for z$\\sim$2 massive galaxies from a combination of UV and 24~$\\mu$m luminosities \\citep{Daddi07a, Wilkins08, Damen09}. {Comparisons between the derivative of the mass density distribution and direct measurements of SFR density also lead to a discrepancy, in particular at z$>$2 \\citep[e.g. ][]{Perez-Gonzalez08}.} Suggestions on how to resolve this discrepancy include changes to the adopted initial mass functions \\citep[IMF, ][]{Dave08} and star formation efficiency \\citep{Khochfar09}, but these cannot be properly tested given the uncertainties associated with the common SFR indicators themselves.  The photo-detector array camera and spectrometer \\citep[PACS, ][]{Poglitsch10} onboard the {\\it Herschel} space observatory \\citep{Pilbratt10} offers a new opportunity to infer the SFR of z$\\sim$2 galaxies from their far infrared (FIR) emission. PACS observes at 160~$\\mu$m (45--65~$\\mu$m rest frame for the sample used here), close to the peak of the emission from dust heated by young stars, away from AGN-heated dust emission and with unprecedented spatial resolution that reduces confusion noise. It allows the most reliable and least biased SFR measurements to date, which avoid large SED extrapolation errors and avoid the uncertainties of attenuation corrections. In this letter, we report on the first measurements of FIR-based calorimetric SFRs of optically selected galaxies at 1.5$<$z$<$2.5 and compare them to results from previous UV and 24~$\\mu$m measurements. We adopt a $\\Omega_m = 0.3$, $\\Omega_\\Lambda = 0.7$, and $H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$ cosmology throughout this letter. ",
        "conclusions": ""
    },
    "1005/1005.1597_arXiv.txt": {
        "abstract": "{ We present {\\it Herschel}-SPIRE observations at 250-500\\,$\\mu$m of the giant elliptical galaxy M86 and examine the distribution of the resolved cold dust emission and its relation with other galactic tracers.  The SPIRE images reveal three dust components: emission from the central region; a dust lane extending north-south; and a bright emission feature $\\rm 10\\,kpc$ to the south-east. We estimate that $\\rm \\sim 10^6\\,M_\\odot$ of dust is spatially coincident with atomic and ionized hydrogen, originating from stripped material from the nearby spiral NGC~4438 due to recent tidal interactions with M86. The gas-to-dust ratio of the cold gas component ranges from $\\sim 20-80$.  We discuss the different heating mechanisms for the dust features.  } ",
        "introduction": "Studies of cold dust in elliptical galaxies has been limited to date by the lack of high-resolution, long wavelength spectral coverage. In particular, the origin of far-infrared (FIR) emission in these systems is still a controversial issue, with evidence of dusty disks favouring a stellar origin (e.g. Knapp et al.\\ 1989) and other systems originating from mergers with dust-rich galaxies (Leeuw et al.\\ 2008).  The unprecedented resolution and sensitivity of the recently launched {\\it Herschel} Space Observatory (Pilbratt et al.\\ 2010) allows us to address long-standing issues such as the origin and quantity of dust in ellipticals.  One of the most well-known IR-bright ellipticals is the giant Virgo Cluster member, M86, at a distance of $\\rm 17 \\,Mpc$ (Mei et al. 2007).  Two dust features were detected with {\\it IRAS}, one coincident with the galaxy and another to the north-west, thought to be coincident with an X-ray plume of gas (Knapp et al.\\ 1989; White et al.\\ 1991) and originally attributed to dust stripped from M86 due to its motion through the cluster.  Higher resolution data from {\\it ISO} revealed two dust peaks within M86 suggesting a massive dust complex (Stickel et al.\\ 2003, hereafter S03).  They proposed a tidal origin, also supported by absorption features attributed to dust stripped from the nearby dwarf galaxy, VCC~882 (Elmegreen et al.\\ 2000).  The discovery of atomic gas offset from the centre of M86 and decoupled from its stellar disk supports the tidal scenario (Li \\& van Gorkom 2001).  More recently, Kenney et al.\\ (2008, hereafter K08) detected strong $\\rm H{\\alpha}$ features extending from the nearby spiral NGC~4438 at $23\\arcmin$ away, to within $1\\arcmin$ of the nucleus of M86.  The distribution and velocity of the ionized gas provides clear evidence for tidal interaction between these two giants (K08, Fig.~1 Cortese et al.\\ 2010).  In this scenario, we are observing debris from the spiral left in the wake of the collision with M86 with $\\sim 10^9\\,\\rm M_\\odot$ of cold gas removed from NGC~4438 (K08).  The stripped material is then heated by the hot interstellar medium (ISM) of M86 or shock fronts from the interaction.  Although tidal stripping from NGC~4438 is supported by the atomic and ionized gas distribution, the origin of the dust responsible for the FIR emission and its heating mechanism is unclear.  Here we present submm observations of M86 with {\\it Herschel}-SPIRE (Griffin et al. 2010). Companion observations of NGC~4438, are presented in Cortese et al.\\ (2010). ",
        "conclusions": ""
    },
    "1005/1005.1205_arXiv.txt": {
        "abstract": "The concept of kinematic assemblages evolving from dispersed stellar clusters has remained contentious since Eggen's initial formulation of moving groups in the 1960's.  With high quality parallaxes from the Hipparcos space astrometry mission, distance measurements for thousands of nearby, seemingly isolated stars are currently available.  With these distances, a high resolution spectroscopic abundance analysis can be brought to bear on the alleged members of these moving groups.  If a structure is a relic of an open cluster, the members can be expected to be monolithic in age and abundance inasmuch as homogeneity is observed in young open clusters.  In this work we have examined 34 putative members of the proposed Wolf 630 moving group using high resolution stellar spectroscopy. The stars of the sample have been chemically tagged to determine abundance homogeneity and confirm the existence of a homogeneous subsample of 19 stars.  Fitting the homogeneous subsample with Yale-Yonsei isochrones yields a single evolutionary sequence of $\\sim$2.7 $\\pm$ 0.5 Gyr.  It is concluded that this 19 star subsample of the Wolf 630 moving group sample of 34 stars could represent a dispersed cluster with an $<$[Fe/H]$>$=-0.01 $\\pm$ 0.02 and an age of 2.7 $\\pm$ 0.5 Gyr.  In addition, chemical abundances of Na and Al in giants are examined for indications of enhancements as observed in field giants of old open clusters, overexcitation/ionization effects are explored in the cooler dwarfs of the sample and oxygen is derived from the infrared triplet and the forbidden line at $\\lambda$6300 \\AA. ",
        "introduction": "A major goal of modern astronomy is to piece together the dynamic and chemical evolution of the Galactic disk.  To this end, one of the principle approaches for probing the disk has been to study open clusters.  Clusters are valuable astrophysical tools as they share common distances, common ages and common initial chemical abundances. With the disk richly populated by both field stars and open clusters, and considering that clusters are relatively well studied, the logical step in piecing together a more complete picture of the chemical and dynamical history of the disk is to study field stars.  In recent years, the advent of large surveys such as \\emph{HIPPARCOS} \\citep{1997A&A...323L..49P} has yielded precise parallaxes for thousands of nearby field stars, and in doing so, provided the necessary tools for investigating the field.  In particular, studies of the velocity distributions of disk field stars in the solar neighborhood have identified stellar overdensities in kinematic phase space (\\citet{1999MNRAS.308..731S}).  The potential application of these velocity structures, commonly referred to as moving groups, was first identified by \\citet{1958MNRAS.118...65E} who considered these assemblages to be relic structures of dissolved open clusters. In this paradigm, a moving group is essentially a spatially unassociated open cluster; therefore it should possess some of the same characteristics that make open clusters such valuable astrophysical tools (common ages and common initial chemical abundances) and similar techniques that are useful for studying open clusters could be applied.  Relatively little work has been done to explore the reality of smaller moving groups (kinematic assemblages of $\\sim$ 100 stars) as dissolved open clusters and their use in chemically tagging the galactic disk, with two notable exceptions: the Ursa Major Group and the HR1614 Moving Group.  \\citet{1993AJ....105..226S} examined the Ursa Major moving group and utilized age information inferred from chromospheric emission to constrain group membership in UMa.  While this study did not utilize chemical tagging to constrain group membership, it did illustrate the viability of moving groups as dissolved populations of open clusters.  \\citet{2003AJ....125.1980K} and \\citet{2005PASP..117..911K} revisited the membership of the UMa group, using new and extant abundances.  They used the results to constrain membership in the UMa group, showed the members to be chemically homogeneous, and noticed overexcitation/overionization effects in the cooler field star members of the group, similar to those observed in young ($<$ 500 Myr) cool open cluster dwarfs (\\citet{2003AJ....125.2085S}, \\citet{2004ApJ...602L.117S}).  The first in depth application of chemical tagging to constrain moving group membership was by \\citet{2007AJ....133..694D}, who derived abundances for various elements for the HR 1614 moving group.  They found that for their 18 star sample, 14 stars were metal-rich ([Fe/H] $\\geq$ 0.25 dex with $\\sigma$=0.03) leading to the conclusion that the HR 1614 moving group, with its distinct kinematics and distinctly super-solar chemical abundances, was a remnant of a dissolved open cluster.  In the field of moving group populations, the classical Wolf 630 moving group is an intriguing target. The first identification of the Wolf 630 moving group was made by \\citet{1965Obs....85..191E} who noted that several K and M dwarfs and giants in the solar neighborhood appeared to have similar space motions to that of the multiple star system Wolf 630 ((U,V,W)=(23, -33, 18) kms$^{-1}$).  These kinematics, distinctive of membership in an old disk population, placed the stars in a relatively sparsely populated region of kinematic phase space \\citep{1969PASP...81..553E}.  Eggen also noted that the color magnitude diagram for the K and M dwarfs and giants with kinematics similar to those of Wolf 630 appeared to trace an evolutionary sequence similar to the old ($\\sim$ 5 Gyr; \\citet{1999AJ....117..330J}) M67 open cluster .  Although his sources are not completely transparent, at least some (17 of 54 stars) of the distances in his study were determined from trigonometric parallaxes , with the remainder coming from ``luminosity estimates of many kinds''. As a rudimentary form of chemical tagging, \\citet{1970PASP...82...99E} estimated metallicities of 23 Wolf 630 group members through \\emph{uvby-$\\beta$} photometry.  Variations in the $\\delta$[m$_1$] index were found to be comparable to the Hyades, Praesepe, and the Coma Berenices clusters, implying chemical homogeneity.  \\citet{1979LIACo..22..355T} studied the chemical composition of five field giant stars that were alleged members of Wolf 630 using high dispersion coud\\'{e} spectra described in \\citet{1979LIACo..22..355T}. Employing a curve of growth approach and measured equivalent widths, they found that three stars appeared to be chemically homogeneous with an overall metallicity for Wolf 630 of [Fe/H]$\\sim$+0.23.  However, it must be noted that their abundances were not measured with respect to the Sun, but are instead quoted with respect to a standard star of presumed solar metallicity (HD 197989), which has since been determined to be a K0III.  While they derived a metallicity of 0.00 for their reference star, literature determinations suggest a value of -0.24. This would lower the average metallicity for the group to [Fe/H]$\\sim$ -0.02.  \\citet{1983MNRAS.204..841M} revisited the membership of the Wolf 630 moving group by recreating the approach presumably utilized by \\citet{1965Obs....85..191E} to find his original Wolf sample.  In summary, they calculate the parallax that yields a V velocity for each group candidate equal to the assumed group velocity of V=-32.8 $\\pm$ 1.3 kms$^{-1}$. The final absolute magnitudes they report assume these parallaxes.  Typical uncertainties in their absolute magnitudes appear to be between 0.2-0.4 magnitudes, larger than magnitude uncertainties obtainable with precise parallax information currently available from Hipparcos.  The color-magnitude diagram assuming these M$_{V}$ values was compared to the scatter of apparent members with the observed scatter in the old open cluster M67.  They concluded that either (1) the intrinsic scatter in the Wolf 630 moving group color-magnitude diagram was greater than that of M67, or (2) the errors in radial velocities and/or proper motions they utilized must have been underestimated by a factor of 2.4 or (3) many of the stars in their sample were, in fact, non-members.  \\citet{1994AAS...185.4516T} examined metallicities from ``published values of [Fe/H] from diverse papers'' of 40 members of the Wolf 630 group.  His sample contains 26 \\% of Eggen's original objects \\citep{1969PASP...81..553E}.  He concluded that metallicity dispersions within his sample were too great for meaningful conclusions about the existence or non-existence of a genuine, chemically distinct Wolf 630 moving group.  This suggests the need to obtain high quality [Fe/H] determinations with minimal uncertainties in testing for chemical uniqueness in a putative Wolf 630 sample.  The analysis of solar neighborhood Hipparcos data by \\citet{1999MNRAS.308..731S} indicates a kinematic rediscovery of the Wolf 630 group.  Their figure 10, showing the \\emph{UV} velocity distribution for 3561 late type dwarfs in the solar neighborhood presents a clear overdensity of stars near the position of Wolf 630.  Furthermore, this structure appears to be distinctly separated from any other known moving groups or stellar streams.  This provides compelling evidence that Wolf 630 is a real kinematic structure.  The question to be asked is if this kinematic structure is composed of stars with a common origin?  Despite the distinctive kinematics exhibited by the Wolf 630 moving group when examined with updated Hipparcos parallaxes, it has not been specifically targeted in an abundance study which makes use of the modern astrometric and spectroscopic data.  This is remedied in this paper, where accurate parallaxes and photometry from the updated \\emph{HIPPARCOS} data reduction \\citep{2007A&A...474..653V} coupled with high precision radial velocities from CORAVEL (\\citet{2004A&A...418..989N} and references therein) allow for developing a Wolf 630 sample with internally consistent distances and absolute magnitudes, thereby removing the uncertainties faced by \\citet{1983MNRAS.204..841M}.  Furthermore, our uniform high resolution spectroscopic study of Wolf 630 moving group candidate members provides a single, consistent set of metallicites with low internal uncertainty to test chemical homogeneity in the group, removing the largest source of uncertainty from \\citet{1994AAS...185.4516T}. ",
        "conclusions": "The primary goal of the paper is to determine if the kinematically defined Wolf 630 Moving Group represents a stellar population of a single age and chemical composition.  The sample stars have been plotted in the HR diagram (Figure \\ref{HR_spec_final}) to determine if they are coincident with a single evolutionary sequence.  The sequence traced by the majority of stars coincides with a Yale-Yonsei isochrone \\citep{2004ApJS..155..667D} of 2.7 $\\pm$ 0.5 Gyr with an assumed solar metallicity. In attempting to qualitatively use ages as a constraint for establishing membership in a distinct evolutionary sequence, it will be assumed that the isochrone which fits the majority of the sample provides a reasonable estimate of the age range of a dominant coeval group, if it indeed exists.  The abundance results are presented in Table \\ref{abundances} and as plots of [X/H] versus temperature (Appendix). Lithium and oxygen abundances were also derived, but they are presented and discussed separately as the approach utilized for these abundance results involved synthesis (Li) or use of the MOOG \\emph{blends} driver (O).  In order to visually present the abundance results, the metallicity distribution of the entire sample is presented in the form of a ``smoothed histogram'' in Figure \\ref{fe_smooth}. This distribution has been generated by characterizing  each star with a gaussian centered on its mean [Fe/H] with standard deviation equal to the [Fe/H] uncertainty. The distributions are summed to yield a final smoothed histogram and have been renormalized to a unit area.  In this manner, the distributions include uncertainties in abundances, making them useful for a visual examination of the complete sample to discern if any stars yield abundances that deviate from the sample as a whole.  The distribution is clearly not unimodal or symmetric.  It is dominated by a near-solar metallicity peak and two smaller peaks at [Fe/H]$\\sim$-0.50 and [Fe/H]$\\sim$+0.30.  It is clear that our Wolf 630 moving group sample is not characterized by a single chemical composition.  \\subsection{Approach to Chemically Tagging}  While our entire sample cannot be characterized by a single chemical abundance, we can investigate whether there is a dominant subsample having common abundances and age. This is done by eliminating stars that are clearly outliers, using arguments based on extreme abundances, evolutionary state (inferred from HR diagram positions, lithium abundance, chromospheric acitivities and surface gravities) or a combination thereof. These members will be classified as ``unlikely'' members of a dominant homogeneous group.  In this way we can, for example, establish a subsample that is characterized by a dominant [Fe/H], if it exists. Stars with such an [Fe/H] will be classified as either ``possible'' or ``likely'' members of a chemically homogeneous, isochronal population having common kinematics.  The final distinctions between ``possible'' and ``likely'' will be made based on evolutionary status and additional abundance information inferred from lithium, alpha elements and iron peak elements.  Particular interest is paid to the iron abundance, [Fe/H], as it is considered the most well determined abundance, primarily due to the quality and size of the Fe line sample.  The quantitative constraint adopted for determining chemical homogeneity was to require that a star's abundance, within its uncertainty, rest within a metallicity band centered on the weighted mean abundance of stars in the sample.  The half-width of this band was conservatively taken to be 3 times the uncertainty in the weighted mean.  This approach was followed in an iterative fashion where whenever a star was determined to be an ``unlikely'' member of a dominant chemical group it was removed from the sample and a new weighted mean and band size was found.  In this manner, a common abundance for the sample was converged to for each element (except Lithium and Oxygen). Examples of the band plots for [Fe/H] versus T$_{\\rm eff}$ is given in Figure \\ref{fe_abtemp}, where [Fe/H] is plotted versus temperature.  The solid line gives the weighted mean [Fe/H] while the dotted lines give the 3-$\\sigma$ uncertainties in this mean, i.e. the abundance band.  This visual analysis from examining the abundance distributions served as a guide for identifying the clearly unlikely members.  Abundance information alone was used to constrain giant star membership in a dominant chemical group, as robust discriminants of age are unavailable.  Many of the dwarfs lay above the main sequence, leading to the question of if they might be pre-or-post main sequence objects. Consequently a diagnostic was needed to constrain evolutionary status for these dwarf and subgiant stars. The full analysis, therefore, examined each star individually, utilizing abundances and information on evolutionary status (inferred from chromospheric activities, isochrone ages and surface gravities) to classify each star in its appropriate category (unlikely, possible or likely).  Figure \\ref{lithium} shows the absolute lithium abundance versus effective temperature for the litle-evolved stars in our sample and for a sample of dwarf stars in the Pleiades, Hyades, NGC752 and M67.  The lithium abundances of the sample stars are plotted with each cluster: filled hexagons are dwarfs, filled triangles are upper limits for dwarfs, open hexagons are subgiants (as inferred from HR-diagram positions and apparently low levels of chromospheric activity) and open triangles are upper limits for subgiants.  Accepted ages are given for each of the respective clusters, with the Pleiades trend being used as a baseline to indicate that a star is likely to be young (i.e. if a star has a lithium abundance which rests in the Pleiades lithium abundance trend it is likely a young star).   \\subsection{Final Membership}  With the considerations above, the 34 stars in the sample have been classified as unlikely, possible and likely members of a common chemical, temporal and kinematic assemblage.  There were a total of 13 stars removed from group membership due to classification as unlikely members. If the remaining 21 stars classified as possible and likely are considered to represent a chemically distinct group, then out of the original kinematically defined sample, $\\sim$ 60\\% remain members of a kinematically and chemically related group with a common 2-3 Gyr age insofar as we can tell.  The final evolutionary sequence traced by the possible and likely members is presented in Figure \\ref{HR_spec_final}, with possible members plotted in red and likely members plotted in green.  The group is reasonably well traced by an evolutionary sequence of $\\sim$ 2.7 Gyr (solid line) with lower and upper limits of 2.2 Gyr and 3.2 Gyr (dashed lines).   The dwarf members, HIP 41484, HIP 105341, HIP 14501 and HIP 43557, have positions that place them slightly above the main sequence; however, based on lithium abundances, none of the stars are believed to be pre-main sequence objects and surface gravities are all consistent with dwarf status. The giants HIP 3992, HIP 34440 and HIP3455 appear to form a red giant clump.  The remaining members all lay on the best fit isochrone within their respective uncertainties.  Thus the possible and likely members we identify can be characterized by a distinct evolutionary sequence of 2.7 $\\pm$ 0.5 Gyrs.  The final UV kinematic phase space plot is presented in Figure \\ref{uvw_final}, where possible members are again red and likely members are green.  For our initial full sample, the RMS U and V velocities are 23.92 and 34.46 kms$^{-1}$, respectively.  In the final subsample of group members, U$_{RMS}$=25.21 kms$^{-1}$ and V$_{RMS}$=35.8 kms$^{-1}$, therefore the kinematic identity has not been significantly altered by the requirement of chemical and temporal coherence to establish group membership, which points to the necessity to utilize criteria other than kinematics to robustely link members of moving groups.  The weighted mean abundances of the final possible and likely members of a dominant chemical group are presented in Table \\ref{group}.  The quoted errors are the uncertainties in the weighted mean.  In order to explore the homogeneity of our samples a reduced chi-squared statistic is presented for each element assuming a constant mean abundance. Performing this test for [Fe/H] for warm stars (T$\\ge$ 5000 K) in the Hyades cluster sample data from \\citet{2006AJ....131.1057S}, yields a $ \\chi_{\\nu}^{2}$ of 1.303. For a set of 7 Pleaides stars from \\citet{2003AJ....125.2085S}, the reduced chi-squared in [Fe/H] is 1.818.  Note that the cool stars were removed from the calculation as they are believed to be impacted by overexcitation/ionization effects. From these chi-squared values, we estimate the Hyades and Pleiades are chemically homogeneous with a roughly 2-sigma significance.  With these open clusters assumed to be chemically homogeneous, an approximate reduced chi-squared of $\\le$ 2, therefore, provides a rough quantitative indication of homogeneity. The $ \\chi_{\\nu}^{2}$ is presented for the full sample of 34 stars ($\\chi_{\\nu}^2all$), the final sample of 21 possible and likely group members ($\\chi_{\\nu}^2group$) and the 11 likely members ($\\chi_{\\nu}^2likely$). First, the very large $\\chi_{nu}^{2}$ for the full sample confirms that the initial kinematically defined sample of alleged Wolf 630 members is clearly not chemically monolithic. The decrease in reduced chi-squared between the full sample and the chemically distinct subsample demonstrates that chemically discrepant stars have been removed.  Even in the likely subsample the $\\chi_{\\nu}^2$ values remain uncomfortably large for Na and Al. Discussion of these discrepancies is reserved for a later section.  Considering the reduced chi-squared for other homogeneous open cluster samples is comparable to the reduced chi-squared for the possible and likely members of the sample across multiple elements, the chosen sample is considered to represent a chemically consistent group with a weighted average metallicity of [Fe/H]=-0.01 $\\pm$ 0.02 (uncertainty in the weighted mean). {\\textbf Using precise chemical tagging of the 34 star sample of the Wolf moving group, a single evolutionary sequence of 2.7 $\\pm$ 0.5 Gyr and [Fe/H]=-0.01 $\\pm$ 0.02 has been identified for a subsample of 19 stars.}   \\subsection{Open Clusters and Moving Groups: Chemically Tagging the Disk}  We present additional results here that illustrate the application of moving group field star members in exploring stellar and chemical evolution in the Galactic disk.  \\subsubsection{Na and Al Abundances}  The abundances of Na and Al appear to be enhanced for some of the stars in the sample. Similar enhancements have been observed in many open clusters.  Most recently, an analysis of abundances in the Hyades cluster found abundance enhancements in Na and Al of 0.2-0.5 dex in giant stars when compared with dwarfs (\\citet{2009ApJ...701..837S}) in line with observations of giant stars in old open clusters (\\citet{2005AJ....129.2725F}, \\citet{2008AJ....135.2341J}). These enhancements can be compared to those observed in the group members of this work.  Plots of [Na/Fe] (top panel) and [Al/Fe] (bottom panel) versus surface gravity are presented in Figure \\ref{na_al}. For the members of the group, the Na and Al enhancements are relatively modest, as seen in a relatively slight upward shift in abundances between dwarfs and subgiants.  The giant abundances, in general, can be brought into agreement with dwarf abundances with downward revisions of 0.1-0.2 dex, consistent with NLTE corrections found in field clump giants with surface gravities down to log $g$=2.10 (\\citet{2006A&A...456.1109M}).  The single star which has greatly enhanced [Na/Fe] and [Al/Fe], HIP 114155, is an evolved, metal poor red giant with enrichments of 0.53 dex and 0.51 dex, comparable to those found by \\citet{2009ApJ...701..837S}. According to the NLTE correction table of \\citet{2003ChJAA...3..316T}, the recommended NLTE correction is at most -0.10, although the calculations performed do not extend below a temperature of 4500 K.  \\citet{1999A&A...350..955G} performed an extensive set of NLTE corrections for Na, and based on their results, there is a recommended NLTE correction of $\\sim$ 0.20 dex.  Even considering these corrections, the Na abundance remains enhanced.  Although there are few NLTE corrections for Al in the literature, \\citet{2008A&A...481..481A} suggest NLTE corrections of roughly 0.60 dex upward.  This is opposite to the necessary correction to remove the enhancement, however the corrections are for low-metallicities ([Fe/H]$\\approx$-2.00).  Further NLTE calculations for cool, moderately low metallicity giant like HIP114155 are needed to determine whether the enhanced abundances in this star are a result of NLTE effects.  The other points of interest in Figure \\ref{na_al} are the two dwarfs with the greatest surface gravities ([Na/Fe]=-0.38 in HIP105341 and [Na/Fe]=-0.33 in HIP5027).  Closer inspection shows that these are the two coolest dwarfs in the sample, perhaps pointing to overexcitation/ionization as a culprit for decreased abundances, similar to overexcitation/ionization effects observed in cool open cluster dwarfs (\\citet{2003AJ....125.2085S}, \\citet{2004ApJ...603..697Y}, \\citet{2005PASP..117..911K} and \\citet{2006AJ....131.1057S}).  Similar effects are not apparent for [Al/Fe].  A single Na line was measurable with a relatively low excitation potential of 2.10 eV, while two Al lines of 3.14 eV and 4.02 eV were used. Additionally, the ionization potential of Al is $\\sim$ 0.9 eV higher than for Na.  These differences are qualitatively consistent with those needed for overexcitation/overionization to be manifest. This can be further explored by comparing abundances from \\ion{Fe}{1} and \\ion{Fe}{1}I.  \\subsubsection{Overexcitation and Overionization in Cool Dwarfs: \\ion{Fe}{1} and \\ion{Fe}{1}I Abundances}  In order to more closely examine the possible effects of overexcitation and overionization for the sample, abundances have been derived from \\ion{Fe}{1} and \\ion{Fe}{2} lines using physical surface gravities (spectroscopic gravities are unsuitable for this purpose since ionization balance forces agreement between abundances of \\ion{Fe}{1} and \\ion{Fe}{2}). Refer to Figure \\ref{overionization_dwarfs} where the difference in abundances between ionized and neutral Fe are plotted versus temperature.  For stars warmer than 4500 K the general trend reveals no overionization within the uncertainties.  The same two coolest dwarfs which evince unusually low [Na/Fe], show large degrees of Fe overionization,  Source of overionization in cool dwarfs are not well-understood, however, one possible explanation is that the stars are active young dwarfs and, thus, heavily spotted.  Recent work suggests that heavily spotted stars have radii which are ``puffed'' compared to standard stellar models (\\citet{2002ApJ...567.1140T}, \\citet{2008A&A...478..507M}. An increased radius would decrease the surface gravity of the star compared to unspotted analogs, which would result in increased \\ion{Fe}{2} line strengths via overionization.  In order to explore the viability of this explanation, the radius that corresponds to the surface gravity needed to eliminate the abundance difference between [\\ion{Fe}{2}/H] and [\\ion{Fe}{1}/H] was determined for HIP5027.  A surface gravity of 3.57 was found to produce agreement between abundances from \\ion{Fe}{1} and \\ion{Fe}{2}, holding temperature and microturbulence constant.  From Yale-Yonsei isochrones, a mass of 0.66 M$_{\\odot}$ is assumed.  The radius for this gravity is R=2.19 R$_{\\odot}$. The radius corresponding to this mass and the physical surface gravity of log $g$=4.70 is R=0.60 R$_{\\odot}$.  From \\citet{2008A&A...478..507M} an upper limit that can be expected for radius changes in this ``spotted'' regime is $\\sim$10\\%, well beneath the radius change implied by the necessary surface gravity change to eliminate overionization and well outside of the uncertainty in the physical surface gravity.  This points to a more likely scenario of significant NLTE effects yielding increased overionization as a function of decreasing temperature as observed in many cool open cluster dwarfs (\\citet{2004ApJ...603..697Y}, \\citet{2006AJ....131.1057S}).   \\subsubsection{Oxygen Abundances: Moving Groups Versus Open Clusters}  Abundances for the $\\lambda$7771, $\\lambda$7774, $\\lambda$7775 high excitation potential oxygen triplet have been derived from equivalent widths. Since abundances derived from the triplet are believed to be enhanced by NLTE effects, corrections from the work of \\citet{2003A&A...402..343T} have been applied to derive NLTE corrected abundances from the triplet lines. The equivalent widths for the triplet, the LTE oxygen abundances, and the final NLTE oxygen abundances are shown in Table \\ref{oxygen}.  The $\\lambda$7774 {\\AA} and $\\lambda$7775 {\\AA} lines appear to be enhanced as a general function of decreasing temperature in both dwarfs (Figure \\ref{oxygen_dwarf_diff}) and giants.  A similar enhancement of the central line (7774.1 {\\AA}) in Hyades giants was noted by \\citet{2006AJ....131.1057S}.  They believed this enhancement to be due to a possible blend with an \\ion{Fe}{1} feature at 7774.00 {\\AA}. While the nature of any blending for the reddest feature (7775 \\AA) is unclear, visual inspection of the spectral line reveals a slight asymmetry, possibly indicating a blend. The distinct increase in [O/H] abundances derived from the red features of the triplet as a function of decreasing temperature suggest that only the blue line (7771.1 {\\AA}) of the triplet should be used for oxygen abundance determinations in cooler stars.  In order to test the possibility of an Fe blend as discussed above, two cool stars of the sample with no measurable oxygen abundances (HIP5027 and HIP105341) were examined to see if they showed any indications of an Fe blending feature near 7774 {\\AA}. In HIP5027 a possible detection of a feature at 7774 {\\AA} was found to have a measured equivalent width of roughly 6.0 m{\\AA}.  This strength is not inconsistent with the expected contribution required from two nearby \\ion{Fe}{1} features at 7773.979 {\\AA} and 7774.06 {\\AA} for the derived Fe abundance.  Neglecting the two red triplet lines in the cool dwarfs, the [O/H] trend of the our dwarf sample is plotted along with the Pleiades trend from \\citet{2004ApJ...602L.117S} (where [Fe/H]=0.00 was assumed to calculate [O/Fe]), and the Hyades trend of \\citet{2006AJ....131.1057S} (where [Fe/H]=$+$0.13 was assumed to calculate [O/Fe]) in Figure \\ref{oxygen_dwarf_simon}. Using $\\lambda$7772 triplet-based [O/H] abundances in 45 Hyades dwarfs, they found a remarkable increase in [O/H] as a function of decreasing temperature for stars with T$_{eff}$$\\le$5400 K.  The increase of [O/H] in the $\\sim$ 120 Myr old Pleiades appeared to be steeper than that in the $\\sim$ 625 Myr old Hyades, perhaps pointing to an age-related effect whereby [O/H] enhancements in cooler stars decrease as a function of increasing age.  Our field dwarfs do not show a drastic increase in abundance as a function of decreasing temperature. The single star that appears to reside within the increasing Hyades trend at cooler temperatures is metal weak (HIP 42499, [Fe/H]=-0.56), resulting in [O/Fe]=+0.47.  The enhanced [O/Fe] ratio at this low metallicity is unsurprising and coincides with the characteristic field dwarf enhancements observed as a function of decreasing temperature for oxygen in other metal poor field stars \\citep{1989ApJ...347..186A}. If the abundance trend observed by \\citet{2004ApJ...602L.117S} and \\citet{2006AJ....131.1057S} is age dependent, the lack of a distinct trend of increasing [O/Fe] with decreasing abundance may point to the stars in the sample being older than the Hyades, not inconsistent with the 2.7 Gyr age of the dominant subsample identified above. If not an age-related effect, then an as yet unknown dichotomy between oxygen abundances in field stars and cluster stars would have to be explored with abundances of field stars of quantifiable age.  For the giant stars in the sample, oxygen abundances have been derived from the infrared triplet and from the forbidden line at $\\lambda$6300.301 {\\AA}, through use of the \\emph{blends} driver of MOOG, following the approach of \\citet{2006AJ....131.1057S}.  In examining the giant triplet abundances, a similar effect as in the dwarfs is observed as temperatures decrease with enhancements in oxygen abundances derived from both the 7774 {\\AA} and 7775 {\\AA} lines. NLTE corrections were applied to the $\\lambda$7771 triplet abundances by interpolating within the grids of \\citet{2003A&A...402..343T}.  The results of these corrections are presented in Figure \\ref{giant_oxygen_forbidden} where forbidden minus permitted [O/H] differences versus temperature are plotted.  Notice that as the temperature decreases, the abundance from the redder lines of the triplet appear to be enhanced relative to the forbidden line.  While the NLTE corrections decreased the abundance enhancements in the cooler stars of the sample, they did not eliminate them.  This yields further evidence of blending effects in the reddest lines of the triplet as a function of coolier temperature. For the purposes of this paper, the oxygen abundances derived from the red features of the triplet will not be used.  In Figure \\ref{giant_forbid_bluetriple} the difference in abundance from the forbidden oxygen line (6300.34 {\\AA}) and the NLTE corrected blue triplet line (7771 {\\AA}) is plotted versus temperature (top plot) and surface gravity (bottom plot).  The dotted line shows a zero difference between the two abundances.  The NLTE-corrected permitted oxygen abundances (7771 {\\AA}) appear to agree well with the forbidden oxygen abundance (6300 {\\AA}) indicating that the blue line of the triplet can provide a reliable oxygen abundance when proper care is taken to make the necessary NLTE corrections.  The single outlier is the highly evolved giant HIP 114155.  The larger abundance from the blue triplet feature in this star is believed to be from NLTE effects that are not removed using the corrections of \\citet{2003A&A...402..343T} as the grid for the corrections does not extend below 4500 K.  While the temperature extrapolation is sufficient for less evolved stars (i.e. NLTE triplet abundances in stars with surface gravities above 2.0 all agree with the forbidden abundance, even at temperatures below 4500 K), the corrections for more evolved stars, with surface gravities $\\sim$1.00, are significantly larger. The good agreement between all other forbidden and blue triplet oxygen abundances indicates the inadequacy extrapolating the NLTE corrections in cool, evolved stars.  The final salient point to make regarding the oxygen abundances is to address the alleged CN blending feature previously discussed.  As mentioned, \\citet{1963rspx.book.....D} list CN features at 6300.265 {\\AA} and two features at 6300.482 {\\AA} with gf values of 5.78E-3, 6.82E-3 and 2.01E-3.  Recall that the inability to adequately calibrate a linelist including these features with a high resolution atlas of Arcturus led to the features not being utilized in the derivation of forbidden line oxygen abundances.  With the good agreement between forbidden oxygen neglecting the CN features and the NLTE corrected blue line of the triplet in Figure \\ref{giant_forbid_bluetriple}, it is suggested that the CN blending features may not be important."
    },
    "1005/1005.4420_arXiv.txt": {
        "abstract": "We have numerically integrated the orbits of ejecta from Telesto and Calypso, the two small Trojan companions of Saturn's major satellite Tethys. Ejecta were launched with speeds comparable to or exceeding their parent's escape velocity, consistent with impacts into regolith surfaces. We find that the fates of ejecta fall into several distinct categories, depending on both the speed and direction of launch.  The slowest ejecta follow sub-orbital trajectories and re-impact their source moon in less than one day. Slightly faster debris barely escape their parent's Hill sphere and are confined to tadpole orbits, librating about Tethys' triangular Lagrange points $L_4$ (leading, near Telesto) or $L_5$ (trailing, near Calypso) with nearly the same orbital semi-major axis as Tethys, Telesto, and Calypso. These ejecta too eventually re-impact their source moon, but with a median lifetime of a few dozen years. Those which re-impact within the first ten years or so have lifetimes near integer multiples of 348.6 days (half the tadpole period).  Still faster debris with azimuthal velocity components $\\gta$ 10 m/s enter horseshoe orbits which enclose both $L_4$ and $L_5$ as well as $L_3$, but which avoid Tethys and its Hill sphere. These ejecta impact either Telesto or Calypso at comparable rates, with median lifetimes of several thousand years. However, they cannot reach Tethys itself; only the fastest ejecta, with azimuthal velocites $\\gta$ 40 m/s, achieve ``passing orbits'' which are able to encounter Tethys. Tethys accretes most of these ejecta within several years, but some 1\\% of them are scattered either inward to hit Enceladus or outward to strike Dione, over timescales on the order of a few hundred years. ",
        "introduction": "There is a sense of youthfulness to Saturn's system, as if the system were now being caught in the process of evolving to something quieter. Signs of youth abound, each of which has something to tell us about what happened at Saturn: (i) Saturn's rings appear both by color and dynamics to be tens or hundreds of millions of years old, rather than the billions of years that describe the Solar System. (ii) Enceladus' eruptions are too energetic to have been going on as they are now for more than a small fraction of the age of the Solar System. (iii) The methane in Titan's atmosphere will be chemically destroyed on a time scale of order 10$^8$ yrs, with the hydrogen escaping to space; there is no known source of new methane, nor is there any sign on the surface of the hundreds of meters of photochemical product which would have accumulated had methane been supplied to the atmosphere over the age of the Solar System. (iv) Saturn's satellite system is crowded and has at least three examples of coorbital moons, a phenomenon not seen elsewhere in the Solar System.  One set of coorbital satellites consists of the two small moons Janus and Epimetheus ($\\sim89$ km and $\\sim59$ km in mean radius, respectively), traveling on mutual ``horseshoe'' orbits. The other two sets each consist of a pair of small Trojan companions of larger moons: Tethys is accompanied by Telesto (Smith et al. 1980) and Calypso (Pascu et al. 1980; Seidelmann et al. 1981), while Dione is accompanied by Helene (Lecacheux et al. 1980) and Polydeuces (Porco et al. 2007). Telesto, Calypso, and Helene are all 11-17 km in radius, and were discovered by ground-based observations, while tiny Polydeuces is only 1.3 km in radius and was found with the Cassini orbiter (Porco et al. 2004). From other spectacular Cassini images, Telesto and Calypso appear to be covered smoothly with fine material.  All of these Trojan moons are intriguing because Tethys and Dione are adjacent; furthermore, Tethys is in a 4:2 mixed-inclination-type mean-motion resonance with Mimas, while Dione, Helene, and Polydeuces are all in a 2:1 inner-eccentricity-type mean-motion resonance with Enceladus. The overall impression is paradoxical. Small number statistics precludes surety, but it seems unlikely that chance alone can explain why Dione and Tethys each have two Trojans, while all other major satellites have no such attendants. Rather, there is a sense of underlying order, as if all of this came about in a perfectly ordinary way.  The better to characterize this system, in this study we simulate the orbital evolution of test particles launched from Tethys' coorbital satellites Telesto and Calypso. First we describe the methodology of the numerical integrations and how we prepare the initial conditions. Next we discuss evolution and fates of ejecta from Telesto and Calypso. The results turn out to be rather interesting: Both the speed and direction of ejecta launch are important, with well-defined regimes which correspond to tadpoles, horseshoes, and passing orbits. The key points are best illustrated by strongly reductionist numerical experiments in which ejecta are launched vertically into space at various places along a satellite's equator. We then broaden our compass to consider more realistic distributions of ejecta launched from impact craters. Our results imply that Telesto and Calypso readily exchange ejecta. ",
        "conclusions": "In this paper we have investigated the fates of ejecta from the Trojan companions of Tethys, namely Telesto and Calypso.  In the first numerical experiment, we systematically investigated the fates of ejecta launched from their equatorial regions. The initial conditions are translated and rotated from a topocentric coordinate system (with origin at the Trojan moon's surface whence the ejecta are launched) to a saturnicentric system and fed to the SWIFT numerical integrator (Levison and Duncan, 1994) to propagate the particle trajectories forward in time for 100,000 years. As expected, the lowest velocity ejecta cannot fully escape the source moon's gravity and fall back onto its surface in less than one day, following suborbital trajectories.  Ejecta do not need to achieve full escape speed to go into orbit about Saturn, though; the launch speed $V_0$ only needs to be greater than $V_e^*= \\gamma V_e$ (where $V_e$ is the classical escape speed, and $\\gamma < 1$; see Eq. 1). Particles with $V_0 \\approx V_e^*$ or slightly higher barely escape the Trojan moon's Hill sphere and are confined to tadpole orbits librating about Tethys' $L_4$ point (for ejecta from Telesto, leading) or $L_5$ point (for ejecta from Calypso, trailing). These particles cannot reach Tethys, and instead are eventually re-accreted by Telesto or Calypso respectively. An interesting finding is that the removal times of ejecta which re-impact Telesto or Calypso within the first dozen years or so are discretized to integer multiples of half the tadpole period.  Somewhat faster ejecta with azimuthal velocities $\\gta$ 10 m/s go into horseshoe orbits, which enclose not only the $L_4$ and $L_5$ points but also $L_3$; these particles still cannot reach Tethys or its Hill sphere and are again accreted by either of the Trojan moons. Ejecta from Telesto or Calypso reaching tangential velocities $\\gta$ 40 m/s go into ``passing'' orbits; beyond this point Tethys accretes most particles, but some are strongly scattered inward or outwards towards other saturnian moons in adjacent orbits such as Enceladus or Dione (on timescales from a few to several hundred years). It is interesting to note that the timescale for removal of an ejectum in a horseshoe orbit is usually orders of magnitude longer than for ejecta in either tadpole or passing orbits.  In order to investigate a more realistic distribution of ejecta, we assumed an impact producing a 1-km-radius crater at five ``cardinal'' points of each Trojan moon (apex, subsaturn, antisaturn, north pole, and south pole). The ejection model adopted here is that of Housen et al. (1983), consistent with soft regolith: the range of ejecta speeds is roughly 5 m/s to 30 m/s. As before, initial conditions are integrated for 100,000 years. As per our earlier findings, none of these particles can go into passing orbits; indeed, all of these go into tadpole or horseshoe orbits. If the former, all particles come back to the source Trojan; if the latter, particles can hit either coorbital. However, we noted that most (between 89\\% and 97\\%) of particles come back to the source Trojan moon. Most of the rest go into the other Trojan moon, with a tiny minority ($<$ 1\\%) surviving the 100,000 year integrations. Ejecta from the antapex of motion was launched according to a ``spallation'' model, with higher speeds of 5 m/s to 56 m/s, so a few percent reached Tethys.  The ease with which Telesto and Calypso exchange ejecta suggests that this recycling may help to maintain their similar sizes in the face of impact erosion. We intend to test this hypothesis in future work on the lopsided Helene-Dione-Polydeuces system. A recent preprint (Izidoro {\\it et al.} 2010) finds that a swarm of planetesimals tadpoling with a proto-satellite generally results in the accretion of 1--4 coorbital companions; this appears to be consistent with our conclusions as well."
    },
    "1005/1005.1669_arXiv.txt": {
        "abstract": "We present Very Large Array observations with 80 milliarcsecond resolution ($\\sim 9$ pc) of the recently discovered Galactic-analog \\hto\\/ masers in the Antennae interacting galaxies (NGC~4038/NGC~4039; Arp244). Three regions of water maser emission are detected: two in the ``interaction region'' (IAR) and the third $\\sim 5.6''$ ($\\gtrsim 600$~pc) west of the NGC~4039 nucleus. The isotropic \\hto\\/ maser luminosities range from 1.3 to 7.7 $L_{\\sun}$. All three maser regions are mostly obscured in the optical/near-infrared continuum, and are coincident with massive CO-identified molecular clouds. The \\hto\\/ maser velocities are in excellent agreement with those of the molecular gas. We also present archival VLA 3.6~cm data with $\\sim 0\\farcs28$ ($\\sim 30$~pc) and $\\sim 0\\farcs8$ ($\\sim 90$~pc) resolution toward the maser locations. All three maser regions are coincident with compact 3.6~cm radio continuum emission, and two are dominated by thermal ionized gas, suggesting the presence of natal super star clusters containing the equivalent of a few thousand O stars. We also present detailed comparisons between the radio data and existing HST ACS (optical) and NICMOS (near-IR) data and find that both maser regions in the IAR are also associated with Pa$\\alpha$ emission and neither source is detected shortward of 2 \\mum\\/. These results highlight the potential of using Galactic-analog \\hto\\/ masers to pinpoint sites of young super star cluster formation with exquisite angular resolution. ",
        "introduction": "Water masers are found in the vicinity of $\\sim 70\\%$ of infrared bright (100 \\mum\\/ $>$ 1000 Jy and 60 \\mum\\/ $>$ 100 Jy) ultracompact \\HII\\/ (UCHII) regions in our Galaxy \\citep{churchwell90,kurtz05}, providing excellent signposts for active star formation.  The term ``kilomaser'' was coined to describe extragalactic H$_2$O masers with luminosities comparable to the brightest star formation H$_2$O masers in our Galaxy (e.g.  W49N; $L_{\\rm H_2O}\\sim 1$ L$_\\odot$). Kilomaser isotropic luminosities ($L_{\\rm H_2O} <10$ L$_\\odot$) are much lower than the more widely studied ``megamasers'' found in the nuclear regions of AGN (i.e.  NGC~4258), with luminosities up to 10$^4$ L$_\\odot$ \\citep{barvainis05}. Since ``kilomasers'' can in principle be either of nuclear (i.e. amplifying a background AGN) or star formation origin, we will use the term ``Galactic-analog'' (GA) \\hto\\/ maser to refer to non-nuclear masers with kilomaser luminosity. Natal super star clusters are thought to contain 1000s of UCHII regions, therefore GA-\\hto\\/ masers may help pinpoint sites of young extragalactic cluster formation, though there have been few searches with the required sensitivity. Until recently GA-\\hto\\/ masers had been unambiguously detected toward only six nearby galaxies: LMC, M82, IC~342, IC~10, M33, and NGC~2146 with isotropic luminosities in the range 0.005 to 4 $L_{\\sun}$ \\citep[][and references therein]{whiteoak86,henkel05,castangia08}. The remaining few known \\hto\\/ kilomasers are either associated with AGN activity or their origin is ambiguous.  We recently conducted a Greenbank Telescope\\footnote{The National Radio Astronomy Observatory operates the GBT and VLA and is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}  (GBT) search for \\hto\\/ masers toward four nearby starburst galaxies ($3< D< 22$ Mpc), known to harbor natal super star clusters and without known AGN (or at worst low luminosity AGN) down to a sensitivity level sufficient to detect strong GA-\\hto\\/ masers. Positive detections were found for all four galaxies \\citep[the Antennae galaxies, He2-10, NGC~4214, and NGC~5253;][]{darling2008}, suggesting that GA-\\hto\\/ maser emission may be common in starburst galaxies. The unusual success of this mini-survey arises from its order of magnitude greater sensitivity compared to most previous single dish surveys for either kilomasers or megamasers.  The strongest masers were detected toward the Antennae (D$\\sim22$ Mpc\\footnote{Antennae distance estimates range from $13.8\\pm 1.7$ Mpc \\citep{saviane08} to $22\\pm 3$ Mpc \\citep{schweizer08}; we adopt 22 Mpc.})  interacting galaxies (NGC~4038/NGC~4039; Arp244). Indeed, this violent merger is an ideal laboratory to study how \\hto\\/ maser emission behaves in regions of extreme star formation.  In this paper we report sensitive high angular resolution Very Large Array (VLA) water maser observations of the Antennae galaxies, along with archival 3.6 and 6~cm continuum, CO, and optical/near-IR data in order to pinpoint the maser locations and determine the nature of the emission. ",
        "conclusions": "We have imaged the GBT-discovered GA-\\hto\\/ masers in the Antennae galaxies with exquisite angular resolution ($\\sim 80$ mas), and find two maser regions in the IAR and a third $5\\farcs6$ ($\\gtrsim$ 600 pc) west of the NGC4039 nucleus.  All three masers show excellent kinematic and spatial agreement with dense CO molecular gas. The IAR maser regions are located in areas of high optical/near-IR extinction, and are coincident with thermal ionized gas suggesting the presence of several thousand O$^*$. Indeed, both of these maser regions seem to pinpoint extremely young sites of deeply embedded super star cluster formation. These results highlight the promise of using GA-\\hto\\/ masers to precisely locate the earliest phases of extragalactic cluster formation. The current isotropic luminosity sensitivity limit of 0.6 L$_{\\sun}$, is still brighter than most Galactic \\hto\\/ masers, suggesting we may only be seeing the bright tip of the maser distribution. With its new broad bandwidth correlator, the Expanded Very Large Array will be an ideal instrument for future GA-\\hto\\/ maser studies (the 1.3~cm data presented here could be taken in 1/3 the time, while achieving $10\\times$ better continuum sensitivity)."
    },
    "1005/1005.3493_arXiv.txt": {
        "abstract": "It has recently been shown in high resolution numerical simulations that relativistic collisions of bubbles in the context of a multi-vacua potential may lead to the creation of bubbles in a new vacuum. In this paper, we show that scalar fields with only potential interactions behave like free fields during high-speed collisions; the kick received by them in a collision can be deduced simply by a linear superposition of the bubble wall profiles. This process is equivalent to the scattering of solitons in $1+1$ dimensions. We deduce an expression for the field excursion (shortly after a collision), which is related simply to the field difference between the parent and bubble vacua, i.e. contrary to expectations, the excursion cannot be made arbitrarily large by raising the collision energy. There is however a minimum energy threshold for this excursion to be realized. We verify these predictions using a number of $3+1$ and $1+1$ numerical simulations. A rich phenomenology follows from these collision induced excursions --- they provide a new mechanism for scanning the landscape, they might end/begin inflation, and they might constitute our very own big bang, leaving behind a potentially observable anisotropy. ",
        "introduction": "} \\label{intro}  The recent interest in the phenomenology of the string landscape \\cite{BP} has sparked a resurgence of interests in the physics of bubbles and bubble collisions. An understanding of such collisions promises much: a prediction of observational consequences of such collisions \\cite{AguJoh09,LarLev09,FreKle09,AguJoh08,ChaKle08,ChaKle07,RubSak01}, an insight into the interaction of non-perturbative objects such as domain walls in general, a more complete treatment of the bubble counting measure \\cite{EasLim05,GarSch05}, possibly a new method of scanning the landscape and many other wonderful things. As Coleman \\cite{Col77} succinctly asked in his seminal paper on the fate of the false vacuum ``...bubble walls begin to collide. What happens then? Can such events be accommodated in the history of the early universe?''  The {\\it kinematics} of wall collisions is well understood: provided one knows the nature of the incoming and outgoing walls (if indeed walls are the only by-products), the wall trajectories can be obtained by energy-momentum conservation \\cite{LanMae01,FreHor07, BouFre08a,AguJoh07}. Our goal, on the other hand, is to understand the collision {\\it dynamics}\\cite{AxeKom99}: how does the field configuration evolve through a collision?  One fruitful approach to understanding the non-perturbative physics involved is numerical, namely lattice simulations of bubble wall collisions. This was pioneered almost thirty years ago by Hawking, Moss and Stewart \\cite{HawMos82a}. Despite their relatively (by today's computational standards) low resolution\\footnote{See also \\cite{KosTur92,KosTur91}, where lattice simulations were used to investigate the collisions of sub-horizon bubbles during a first order electroweak phase transition.}, their 1+1 dimensional simulations discovered an interesting result that was not predicted by analytic methods: the collision of two true vacuum bubbles creates a relatively long live pocket of false vacuum which eventually collapse under differential pressure between the spacetime regions. This result was not well understood or exploited until recently.  In high-resolution, 3+1 dimensional, numerical simulations, it was shown that not only is this formation of false vacuum pocket a fairly robust effect, the right kind of potential can produce a new, stable bubble of a different vacuum (\\cite{EasGib09}, henceforth EGHL). Furthermore, such \\emph{classical transitions} can be quite elastic --- coherent bubble walls form immediately after collisions with little dissipative losses in the form of scalar radiation. EGHL also showed that whether such transition occurs is dependent on both the energy of the collision, and the height of the potential barrier between the progenitor bubble vacuum and the progeny bubble vacuum. The setup is illustrated in Fig. \\ref{setup}.  \\begin{figure*}[tb] \\subfigure[] {\\label{setup1}\\includegraphics[width=.50\\textwidth]{Figures/introconf.eps}} \\hspace{0.1in} \\subfigure[] {\\label{setup2}\\includegraphics[width=.40\\textwidth]{Figures/introplot.eps}} \\caption{{\\bf (a).} A space-time diagram showing the collision of two bubbles embedded in a false vacuum sea, in a frame in which the two bubbles nucleate at the same time. We refer to the overlap between the bubbles as the transition region -- here, the scalar field is generally kicked by the collision to somewhere else in field space. {\\bf (b).} A schematic diagram showing the potential for the scalar field. The arrow denotes the scalar field kick as seen by an observer following the black arrow in panel (a). This is the kick immediately after a collision, which is realized as long as a certain minimum condition is met. Where the field will roll to in the long run depends on the precise layout of the landscape in the transition region. Note that only one bubble vacuum is shown. The other bubble might inhabit the same or some other vacuum. Note also the field space does not have to be one-dimensional. } \\label{setup} \\end{figure*}  In this paper, we wish to extend the work of EGHL, beginning with the question: what determines the size of the collision induced excursion? By this we mean the field excursion at $x=0$ (the location of the collision), from immediately before to immediately after the collision (Fig. \\ref{setup}). Naively, one might think that the excursion depends on the collision energy --- the higher the energy, the further the scalar field could go --- as if the kinetic energy stored in the bubble walls can be tapped to initiate motions in field space. This kind of reasoning proves to be misleading, however.  Let us try to gain some insight by simplifying the problem. Recall that bubble walls accelerate as the bubbles grow in size. A high speed collision therefore also means a collision between two large bubbles. In this limit, we can treat the collision event as one between two planar walls, in other words it is a 1+1 dimensional problem. The collision of two bubble walls is thus equivalent to the collision of two kink solitons. Much is known about the scattering of solitons in for instance the famous sine-Gordon model (Appendix \\ref{sinegordon}). We will be able to deduce some general model independent results in the relativistic collision limit, including the case of multiple scalar fields. And we will find that the collision induced field excursion is in fact bounded, and given by a simple expression involving the field difference between the false and the bubble vacua (\\S \\ref{freepassage}). There is a minimum energy threshold that must be reached for this maximal field excursion to be realized, which is derived in \\S \\ref{transition}. These analytic predictions are verified with numerical experiments in \\S \\ref{simulations}. We will digress a bit in \\S \\ref{multiple} to consider cases where the excursion takes the field through multiple local minima.  These collision induced field excursions open up a rich set of questions, including: How do they impact inflation? What are the observational signatures if we reside in the transition region (Fig. \\ref{setup})? What happens when there are multiple collisions? We will explore these questions in \\S \\ref{sect:discussion}. For the most part, we focus in this paper on scalar field models with canonical kinetic terms and potential interactions, and work in flat space. Relaxing these assumptions will be discussed in \\S \\ref{sect:discussion} as well. ",
        "conclusions": "\\label{sect:discussion}  To briefly summarize, we find that linear superposition of wall profiles provides a good approximation to what happens at a high speed collision. The free passage of wall profiles implies a field excursion in the collision region that is given by Eq. (\\ref{multidimkick}). For this free passage kick to be successfully realized, the collision must exceed a certain minimum energy threshold given by Eq. (\\ref{gammacondition}). Beyond this threshold, the size of the kick itself is independent of the collision speed, i.e. one cannot obtain arbitrarily large field excursions by pumping up the collision energy. When the minimum energy condition is met, the free passage kick can cause a transition into a new vacuum, if the kick takes the scalar field to within the new vacuum's basin of attraction. Interesting possibilities arise if the excursion traverses over several new vacua, and they can be understood by similar arguments (\\S \\ref{multiple}). We have verified these statements using numerical computations of soliton collisions in $1+1$ and bubble collisions in $3+1$.  These findings have many interesting theoretical and observational implications. Let us discuss some of them.  \\subsection{Including Gravity} \\label{sect:gravity}  Our simulations used a flat, non-expanding background.  To study a classical transition in the cosmological context, first we need to know the necessary modifications once we include gravitational effects.  Since the collision happens on a short time scale, the corrections from gravity will not be in the collision itself, but in the domain wall motions before and after the collision.  Without gravity, it appears that arbitrarily high barriers can always be traversed via collisions, as long as the incoming boost $\\gamma$ is large enough. In an expanding background, however, there is a maximum boost for every collision, \\begin{equation} \\gamma_{\\rm Max} \\sim (R_0 H_A)^{-1}~, \\end{equation} where $R_0$ is the initial size of the bubble B, and $H_A$ is the hubble expansion rate for the parent vacuum A.  Therefore, it is easy to generalize our results to an expanding background, as long as we maintain an additional condition : the critical $\\gamma$ to make the transition in flat space has to be smaller than $\\gamma_{\\rm Max}$ in the expanding background, otherwise there will be no transition.  After the transition, there are more complications.  For example, if $V_C>V_B$, the domain walls can turn around and may or may not collide with each other again, thereby sealing the vacuum C region.  Also when the domain wall is heavy, $m_{BC} < |V_C-V_B|$, gravity allows it to accelerate and eventually run away from both sides.  These behaviors are hard to keep track of in lattice simulations.  An analytical study on these questions can be found in  \\cite{JohYan10}.  \\subsection{Cosmology After a Collision} \\label{sect:cosmology}  Many recent studies discussed the possibility of observing signals from bubble collisions in our past.  Several cases are summarized in \\cite{AguJoh09}.  One important issue is whether slowroll inflation can be maintained after a violent disturbance like a bubble collision.  The scalar field behavior we study here provides useful intuition to address such a problem.  If the field range of slowroll on the potential is super-planckian, which is usually known as a large-field slowroll inflation, then it is quite stable against disturbances as we do not expect the vacuum separation (which determines the collision induced kick) to be super-planckian. A sufficiently large flat region in the potential post-collision would provide a buffer for slowroll to continue unscathed. On the other hand, if the slowroll range is smaller, then it is more delicate.  One example in \\cite{AguJoh09} seems to suggest that a collision always ends small-field slowroll inflation.  Here we would like to argue the opposite---\\emph{Having small-field slowroll inflation after a collision is as plausible as having it after a Coleman-De-Luccia tunneling.}  We start from Figure 7 in \\cite{AguJoh09}, which the authors interpreted as the collision pushing the field from slowroll directly to the end of inflation. The free passage tells us that right after the domain walls cross each other, the field ends up within a region (possibly very) near the parent vacuum (see Figure 7 of \\cite{AguJoh09}). Because of the pressure difference, such a region undergoes oscillations and re- collisions as in \\cite{HawMos82a}. In the figure cited above, the first two collisions are visible but the rest of them are too small to be resolved, forming an effective domain wall\\footnote{We have reproduced this result in both 3+1 and 1+1.}.  The pressure difference between the bubbles on left and right accelerates this effective domain wall to the right.  The creation of this domain wall is highly dissipative -- scalar radiation is emitted from the point of initial collision and subsequent re-collisions, and this scalar radiation is large enough in amplitude to disrupt the small field inflation, causing the inflaton to topple off its delicately balanced potential and thermalize.  In other words,  the failure of small-field slowroll inflation appears to be a result of the disturbances of the repeated re-collisions hidden in the effective domain wall.  Nevertheless, as we have  shown in this paper, counter to the intuition that collisions are violent, classical transitions can be exceptionally gentle (recall the fairly homogeneous post-collision region in our $1+1$ and $3+1$ simulations). This gentleness might live in harmony with the delicate small-field potential. Consider a potential in Fig.~\\ref{fig:ABCroll}.  If we arrange that $\\phi_A-\\phi_B=\\phi_B-\\phi_C$, then by free passage, a collision between two bubbles of $B$ naturally starts slowroll inflation in $\\phi_C$. \\begin{figure}[tb]\\centerline{\\epsfxsize=8cm\\epsffile{Figures/slowroll.eps}} \\caption{A potential with vacua $A,B$ and a small field slowroll potential around $C$.} \\label{fig:ABCroll} \\end{figure} One might argue that such a setup is finely tuned.  Let us ask a different question: Can a Coleman-De-Luccia tunneling from vacuum $B$ start slowroll inflation in $\\phi_C$?  The answer is yes, if we arrange the potential correctly -- it is also finely tuned!  Although there seems to be an additional parameter $\\gamma$ involved in collisions, we remind ourselves the following two facts. \\begin{itemize} \\item Our simulation shows that once the transition is allowed, increasing $\\gamma$ further has no impact on the field value right after the transition.  Namely, the initial condition for starting the slowroll is not sensitive to the incoming boost. \\item In a multiverse we have an infinite number of collisions with different $\\gamma$'s, so it is naturally scanned and does not require further fine-tuning. \\end{itemize} Therefore, all the fine-tuning is on the potential itself.  We conclude that starting slowroll through collisions is not more fine-tuned than starting it through a single tunneling.\\footnote{These fine-tunings are only necessary in small field inflation models.  In large field models it will be equally generic to start inflation through either tunneling or collision.}  Let us now turn to the question of the spacetime geometry inside a collisionally formed bubble. The motivation is that we might perhaps live in it, i.e. a collision event might have constituted our beginning, our big bang. Assuming that there is slowroll inflation inside the bubble, we can approximately describe the geometry as pure de~Sitter, whose general metric is given in \\cite{FreHor07}: \\begin{equation} ds^2 = -\\frac{dz^2}{f(z)}+f(z)dx^2+z^2dH_2^2~, \\end{equation} where \\begin{equation} f(z) = 1-\\frac{M}{z}+H^2z^2~. \\end{equation}  When the domain walls are highly boosted, the conservation of stress-tensor gives us\\footnote{This is exactly true for massless(null) domain walls \\cite{LanMae01}, and  approximately true for light or highly boosted domain walls.} \\begin{equation} f_A(z_c)f_C(z_c) = [f_B(z_c)]^2~, \\end{equation} where $z_c$ is the collision radius of the two-dimensional hyperbolic section $H_2$.  The $A, B, C$ labels follow the labeling of vacua in the rest of this paper. Neglecting radiation loss (which seems to be a good approximation from our simulations), $M=0$ in both $f_A$ and $f_B$.  As arranged in the potential $V_A>V_B>V_C$, $M$ in $f_C$ is also small and the $M/z$ term soon becomes unimportant.  The metric in region C is thus approximately \\begin{equation} ds^2 = -\\frac{dz^2}{1+H_C^2z^2} + (1+H_C^2z^2)dx^2 + z^2dH_2^2~. \\end{equation}  When there is slowroll inflation in region C, this will not be the proper slicing to describe the universe.  Region C is the forward lightcone of an $H_2$ with radius $z_c$.  Following the usual prescription --- in the case of a single bubble --- for deriving an open de~Sitter slicing covering the future lightcone of a point in the global de Sitter space, we can similarly derive the ``open'' slicing of the universe in region C: \\begin{eqnarray} ds^2 &=& -dt^2 + \\frac{\\sinh^2H_Ct}{H_C^2}d\\xi^2 \\\\ \\nonumber &+& \\bigg(z_c \\cosh H_C t + \\frac{\\sinh H_Ct}{H_C} \\sqrt{1+H_C^2z_c^2}\\cosh\\xi\\bigg)^2dH_2^2 \\end{eqnarray}  We can see that the metric preserves only the symmetry of $H_2$, and is both anisotropic and inhomogeneous.  The interesting trait is that both its anisotropy and inhomogeneity are correlated with its negative spatial curvature. Since the universe is roughly isotropic and homogenous to 1 part to $10^5$ \\cite{WMAP7}, if we are to live in such a bubble, there must exist a sufficient amount of inflation in this new universe after its birth through a collision \\cite{Wal83}. While the possibility of statistical anisotropy have been discussed, at least at the level of the initial perturbations \\cite{AckCar07}, our model predicts an interesting novel anisotropy in the form of an anisotropic spatial curvature, i.e. there exist two different curvature scales in different directions, an intriguing possibility we plan to explore in future work. Note that an anisotropic curvature shifts the Doppler peaks in the microwave background on all angular scales, making the effect easier to observe than most other large scale anomalies, as recently emphasized by \\cite{GraHar10,BlaSal10} in the context of other models.   \\subsection{Additional Ingredients to the Field Dynamics} \\label{sect:addition}  There are several possible complications to the field dynamics beyond what we have discussed. First of all, there can be derivative interactions. Our analysis should hold as long as the collision process does not probe energy scales above the cut-off for these derivative interactions, say $\\Lambda$. This means the boosted thickness of our bubble walls $1/(\\gamma\\mu)$, where $1/\\mu$ is the rest-frame thickness, must be larger than $1/\\Lambda$. This implies an {\\it upper limit} to the collision speed that we can consider $\\gamma \\lsim \\Lambda/\\mu$. When derivative interactions are important, linear superposition such as we have used in Eq. (\\ref{phisum}) no longer works.  Another possibility is that the field metric (in the multi-field case) is non-trivial, and therefore the excursion trajectory is modified. A third possibility is that some other field is coupled to our bubble scalar field, and the coupling is such that this other field becomes massless in the course of an excursion. Massless particles get produced\\footnote{In\\cite{ZhaPia10} it is treated as a direct generalization of preheating, but we remain conservative about their results.}, and could significantly modify the excursion trajectory. We plan to investigate these intriguing possibilities in the future.  \\subsection{Scanning of the Landscape}\\label{sect:landscape}  The collision induced excursion offers a new mechanism for scanning a landscape of many vacua. Our results here touch on two aspects of this scanning. One is that even very high barriers can be overcome by a collision as long as the collision is relativistic enough (but subject to constraints from expansion, and from derivative interactions). It could well be that certain vacua surrounded by high barriers are more likely to be populated by collisions rather than direct tunneling. A second important feature of this scanning is that it is bounded --- the collision induced excursion does not become arbitrarily large by raising the collision speed. Rather, it is controlled by a simple vector sum rule (Fig. \\ref{multiscalarkick}), which tells us that the excursion can only be as large as the field difference between the parent and bubble vacua. In theories in which tunneling between vastly separated vacua is common \\cite{BroDah10}, large collision excursions are also possible. Finally, this excursion has a specific direction, and it is interesting to ask what is the typical basin of attract there.  In\\cite{AguJohLar09} is it  suggested that due to the universal dilatonic runaway direction in string theory inspired models, classical transition can lead to decompactification of extra dimensions.  \\subsection{Multiple Collisions} \\label{sect:multi}  The free passage approximation and the resulting vector sum rule in section \\ref{freepassage} is a simple fact of the field dynamics.  It applies to any solitonic objects interpolating between local minima.  We can arrange a potential in which a collision is followed by further collisions, leading to a nested set of classical transitions. A spacetime structure with multiple, overlapping collisions follows simply from the vector sum rule, as shown in Figure~\\ref{figure:CFT}.  A particularly intriguing version of this occurs when a bubble is nucleated in the presence of a compact dimension \\cite{Bro08}. The bubble grows, wraps around the compact dimension, and eventually collides with itself. A new vacuum opens up between the outgoing pair of walls. These walls eventually collide after traveling through the compact dimension. Further transitions and collisions follow, as far as the potential allows. This provides a novel way to realize an old idea by Abbott \\cite{Abb85}, by classical transitions rather than tunneling.  \\begin{figure}[tb] \\centerline{\\epsfxsize=6cm\\epsffile{Figures/CFT.eps}} \\caption{A spacelike slice through multiple, overlapping collisions. Following the vector sum rule, we have:\\\\ $\\vec \\phi_C=2\\vec \\phi_B-\\vec \\phi_A$, $\\vec \\phi_D=3\\vec \\phi_B-2\\vec \\phi_A$.} \\label{figure:CFT} \\end{figure}  In the context of a scalar field, such a cascade of classical transitions seems to rely on the existence of roughly evenly spaced vacua in the potential landscape. Interestingly, in \\cite{BlaJos09} it is shown that in a model with multiple vacua constructed from extra dimensions and fluxes, a classical transition is the most natural result, as shown in Figure~\\ref{figure:flux}. \\begin{figure*}[tb] \\epsfxsize=5cm\\epsffile{Figures/f6.eps} \\ \\ \\epsfxsize=5cm\\epsffile{Figures/col1.eps} \\ \\ \\epsfxsize=5cm\\epsffile{Figures/col.eps} \\caption{A classical transition (middle) can occur when the positions of vacua are arranged appropriately (left), as demonstrated in this paper.  Alternatively one can have a model like the figure on the right, where vacua 1,2,3 have 1,2,3 units of flux and the domain wall is the charged object which changes the flux.  If the charged objects freely pass through each other, a classical transition is the most natural result of every collision.} \\label{figure:flux} \\end{figure*} In these models, the multiple classical transition structure in Figure~\\ref{figure:CFT} should be taken seriously.  A similar but simpler structure was studied in \\cite{FreKle09a} and suggested the possibility of a CFT description for an eternal inflating spacetime."
    },
    "1005/1005.0490_arXiv.txt": {
        "abstract": "The aim of this paper is to study the latitudinal variation in the solar rotation in soft X-ray corona. The time series bins are formed on different latitude regions of the solar full disk (SFD) images that extend from 80\\degr S to 80\\degr N. These SFD images are obtained with the soft X-ray telescope (SXT) on board the {\\it Yohkoh} solar observatory. The autocorrelation analyses are performed with the time series that track the SXR flux modulations in the solar corona. Then for each year, extending from 1992 to 2001, we obtain the coronal sidereal rotation rate as a function of the latitude. The present analysis from SXR radiation reveals that; (i) the equatorial rotation rate of the corona is comparable to the rotation rate of the photosphere and the chromosphere, (ii) the differential profile with respect to the latitude varies throughout the period of the study; it is more in the year 1999 and least in 1994  and (iii) the equatorial rotation period varies systematically with sunspot numbers and indicates its dependence on the phases of the solar activity cycle. ",
        "introduction": "\\begin{figure*} \\centering{ \\includegraphics[width=\\textwidth]{Figure1.eps}} \\caption{The two panels on the top are typical examples of the time series of average soft X-ray flux of the year 1999 at latitudes 20\\degr south and 20\\degr north. The two panels on the bottom are their autocorrelograms. The peaks of these autocorrelogram are used in the present analysis to determine synodic rotation period.} \\end{figure*}   There have been a series of studies on the coronal rotation rate employing Soft X-Ray (SXR) data \\citep{b19,b14} in the past few decades, yet the picture is far from being completely clear. However, the extensive SXR database acquired by {\\it Yohkoh} Soft X-ray Telescope (SXT) (extending almost the entire solar cycle) contains a wealth of information to be analyzed about the variation of coronal rotation rate with time and latitude. The rotation of the solar corona is still a matter of great interest because of its unique latitudinal and height profiles, temporal and spatial variations and its deviation from rigid rotation to the latitudinal-dependent differential rotation.  Observations of the solar coronal rotation period, in the last decades, have shown variations between rigid rotation to differential rotation profile across latitudes. \\citet{b21} and \\citet{b22} have used the data obtained by the {\\it Yohkoh}/SXT to study the rotation rates of the corona and have reported that the soft X-ray corona does not show any significant differential rotation; it means the corona is more rigid than was thought before, compared to the photosphere or chromosphere.  Recently, by making use of X-ray bright points (XBPs) {\\citet {b13} determined the coronal rotation rate employing solar full disc (SFD) images  obtained through the soft X-ray telescope (SXT) on the {\\it Yohkoh} and the X-ray telescope (XRT) on the {\\it Hinode} solar observatory. {\\citet {b13} found that the corona does rotate differentially throughout the solar magnetic cycle, in the same way as its neighbouring lower atmospheric layers do, but it does not depend on the phases of the solar magnetic cycle.  After analyzing the synoptic photoelectric observations of the green (Fe XIV emission lines at 530.3 nm) and red (Fe X emission line at 637.4 nm) coronal lines for the periods 1973 to 2000 (Fe XIV) and 1984 to 2000 (Fe X), respectively, \\citet{b17} and \\citet{b1,b2} concluded that the rotation of corona shows more compactness than the rotation of the sunspots. The presence of a weak, solar cycle dependent differential rotation in both the spectral lines was also reported in their study. It indicates that the corona shows rigid rotation near the solar minimum and differential rotation near the solar maximum.  On the basis of the analysis of the coronal bright points observed on solar full-disc filtergrams at 19.5 nm (Fe XII line) wavelength with the extreme ultraviolet imaging telescope (EIT) on board the {\\it SoHO} spacecraft, {\\citet {b12}} found that the rotation of coronal bright points closely follows the latitudinal rotation profile of the photospheric magnetic field and, simultaneously, it was also shown that coronal features at different heights in the corona may exhibit different rotation rates. The SFD filtergrams from {\\it SoHO} EIT (Fe XV) instrument at 28.4 nm were also used by {\\citet {b6}} to analyze solar differential rotation by tracing coronal bright points. {\\citet {b6}} determined through the procedure based on {\\it IDL}, that the differential rotation profile corresponds roughly to the rotation of sunspot groups.  \\begin{table*} \\centering \\begin{minipage}{112mm} \\caption{Sidereal rotation periods of soft X-ray corona obtained through autocorrelation analysis for each latitude bins on SFD images for the years 1992-2001.} \\begin{tabular}{@{}lcccccccccc@{}} \\hline Latitude \t& \\multicolumn{10}{c}{Years} \\\\ bins at\t\t&\t1992\t&\t1993\t&\t1994\t&\t1995\t& 1996  &  1997 &\t1998\t&\t1999  &\t2000\t&\t2001\\\\ \\hline 80\\degr S\t\t&\t26.4\t&\t27.7\t&\t25.1\t&\t25.7\t&\t26.0\t&\t26.9\t&\t27.0\t&\t28.0\t& 26.0\t&\t26.4\\\\ 70\\degr S\t\t&\t26.4\t&\t27.1\t&\t25.1\t&\t25.6\t&\t26.0\t& 26.0\t& 27.0\t& 27.7\t& 25.4\t& 26.4\\\\ 60\\degr S\t\t&\t26.5\t&\t26.5\t&\t25.7\t&\t25.6\t& 25.1\t& 26.0\t& 26.9\t& 27.2\t& 25.1\t& 25.1\\\\ 50\\degr S\t\t&\t25.7\t&\t27.2\t&\t25.1\t&\t25.6\t&\t25.6\t& 26.0\t& 26.7\t& 26.0\t& 25.1\t& 24.9\\\\ 40\\degr S\t\t&\t25.6\t&\t25.1\t&\t25.1\t&\t25.7\t&\t25.1\t& 26.0\t& 26.5\t& 26.2\t& 24.3\t& 25.1\\\\ 30\\degr S\t\t&\t26.0\t&\t25.1\t&\t25.3\t&\t25.1\t&\t24.6\t& 26.0\t& 26.0\t& 25.8\t& 24.3\t& 25.6\\\\ 20\\degr S\t\t&\t25.6\t&\t24.6\t&\t25.4\t&\t24.7\t& 24.8\t& 25.1\t& 25.8\t& 25.8\t& 24.7\t& 25.6\\\\ 10\\degr S\t\t&\t24.9\t&\t25.1\t&\t25.1\t&\t24.7\t&\t24.8\t& 25.4\t& 25.7\t& 25.7\t& 24.0\t& 24.3\\\\ Equator\t\t\t&\t24.7\t&\t25.7\t&\t24.8\t&\t25.1\t&\t24.6\t& 24.3\t& 25.0\t& 25.3\t& 24.3\t& 24.3\\\\ 10\\degr N\t\t&\t25.8\t&\t25.6\t&\t24.8\t&\t25.0\t&\t25.1\t& 25.1\t& 25.1\t& 25.1\t& 24.6\t& 23.8\\\\ 20\\degr N\t\t&\t25.7\t&\t25.1\t&\t24.8\t&\t24.8\t&\t24.6\t& 25.1\t& 25.1\t& 24.6\t& 24.9\t& 24.3\\\\ 30\\degr N\t\t&\t26.0\t&\t25.1\t&\t25.0\t&\t25.1\t&\t24.6\t& 25.6\t& 26.0\t& 25.1\t& 24.9\t& 25.4\\\\ 40\\degr N\t\t&\t25.8\t&\t25.6\t&\t25.1\t&\t25.1\t&\t24.6\t& 26.0\t& 26.0\t& 26.4\t& 25.3\t& 25.1\\\\ 50\\degr N\t\t&\t25.6\t&\t26.0\t&\t25.3\t&\t26.0\t&\t25.4\t& 26.0\t& 26.0\t& 26.4\t& 25.5\t& 24.6\\\\ 60\\degr N\t\t&\t25.1\t&\t26.3\t&\t25.7\t&\t24.9\t&\t24.7\t& 26.4\t& 27.3\t& 27.3\t& 25.5\t& 25.3\\\\ 70\\degr N\t\t&\t26.4\t&\t25.1\t&\t25.7\t&\t25.1\t&\t26.0\t& 26.9\t& 27.1\t& 28.2\t& 25.7\t& 26.6\\\\ 80\\degr N\t\t&\t26.4\t&\t26.0\t&\t25.6\t&\t25.7\t&\t25.6\t& 26.9\t& 26.6\t& 28.2\t& 25.4\t& 26.9\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}     Disk integrated solar radio fluxes from 1997 to 1999 at different frequencies are used for the study of coronal rotation at different heights \\citep {b24, b20}. \\citet{b20} have shown that coronal rotation depends on the heights in the corona. Using Nobeyama radio images at 17 GHz for the years 1999 to 2001, \\citet {b7} demonstrated that the rotation profile of the lower corona across the latitude is somewhat shallower than the rotation profile of layers below it.  From a detailed study of the data from {\\it LASCO}, \\citet {b15} concluded that the rotation of the corona displays a radially rigid rotation of 25.57 days sidereal period from 2.5 $R_{\\sun}$  to $>$ 15 $R_{\\sun}$.  The solar corona on the full disc can be best seen in soft X-ray images. This is because, in X-rays radiations, the corona is optically thin but, since the photosphere is comparatively less hot than the corona and therefore does not emit shorter wavelengths, soft X-rays radiations from the solar corona's full disc appear as bright emissions against a dark background of photosphere. Such high contrast images obtained in soft X-rays make it easier to track the coronal features present anywhere on the solar disc and to determine spatial and temporal variations by using the time series SXR data.  The paper is structured as follows. In section 2, we discuss the source and span of the data which we have used in this study. In its subsection, we briefly give the technique used to obtain the time series from the pre-processed solar full disc images. In section 3, we present the method used to determine the rotation period at various latitudes from such time series. In section 4, the rotation rate and its differential profile is compared with recent published results. It also includes the discussion on the dependency of the rotation rate on the activity cycle. The paper ends with conclusions. ",
        "conclusions": "It is clear from Figure 3, that the equatorial sidereal rotation rate we obtained (14.5 deg/day as shown in Table 3) is similar to those obtained by the \\citet{b6} (14.5 deg/day from Table 3) or by the \\citet{b12} (14.9 deg/day from Table 3). Interestingly, they both used as tracers coronal bright points at chromospheric levels ({\\it SoHO}/EIT). A similar inference can be drawn at the photospheric level (14.6 deg/day determined by \\citet{b3} using sunspots as tracer). So from the present study, we infer that the soft X-ray corona at the equator rotates almost at the same rate as the lower atmospheric layers.  The differential profile as a function of latitude varies considerably throughout the period of the study, extending from 1992 to 2001. The differential gradient is found to be maximum in the year 1999 and minimum in 1994. The comparison shows that the curve of the present work using the {\\it Yohkoh} dataset shows latitudinal dependent differential rotation but its mean gradient is less in comparision to most of the previously published results (except that of \\citet{b21}). It implies that the rotation profile of soft X-ray corona across latitudes is shallower than that of chromosphere and photosphere. The result is important in the sense that SXR corona with {\\it Yohkoh}/SXT dataset shows the same qualitative rotational behaviour as that obtained in some other studies, for example, \\citet{b9, b8, b21}, employing different rotation measurement techniques.  The equatorial rotation rate in the solar corona shows a systematic variation with respect to the annual sunspots number. But the variation in its rate is anti-correlated with the sunspots number (except in the year 2000). This indicates the dependence of the coronal equatorial rotation velocity on the phases of the solar cycle."
    },
    "1005/1005.2176_arXiv.txt": {
        "abstract": "We report on the first demonstration of time-delay interferometry (TDI) for LISA, the Laser Interferometer Space Antenna. TDI was implemented in a laboratory experiment designed to mimic the noise couplings that will occur in LISA. TDI suppressed laser frequency noise by approximately $10^9$ and clock phase noise by $6\\times10^4$, recovering the intrinsic displacement noise floor of our laboratory test bed. This removal of laser frequency noise and clock phase noise in post-processing marks the first experimental validation of the LISA measurement scheme. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.0035_arXiv.txt": {
        "abstract": "We report the observation of sixteen cosmic ray events of mean energy of $1.5 \\times 10^{19}$~eV, via radio pulses originating from the interaction of the cosmic ray air shower with the Antarctic geomagnetic field, a process known as geosynchrotron emission. We present the first ultra-wideband, far-field measurements of the radio spectral density of geosynchrotron emission in the range from 300-1000~MHz. The emission is 100\\% linearly polarized in the plane perpendicular to the projected geomagnetic field. Fourteen of our observed events are seen to have a phase-inversion due to reflection of the radio beam off the ice surface, and two additional events are seen directly from above the horizon. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.1613_arXiv.txt": {
        "abstract": "s{ A stochastic background of primordial gravitational waves could be detected soon in the polarization of the CMB and/or with laser interferometers. There are at least three GWB coming from inflation: those produced during inflation and associated with the stretching of space-time modes; those produced at the violent stage of preheating after inflation; and those associated with the self-ordering of Goldstone modes if inflation ends via a global symmetry breaking scenario, like in hybrid inflation. Each GW background has its own characteristic spectrum with specific features. We discuss the prospects for detecting each GWB and distinguishing between them with a very sensitive probe, the local B-mode of CMB polarization. }    Cosmological Inflation~\\cite{LindeBook,MukhanovBook} naturally generates a spectrum of density fluctuations responsible for large scale structure formation which is consistent with the observed CMB anisotropies.\\cite{Komatsu2010} It also generates a spectrum of gravitational waves, whose amplitude is directly related to the energy scale during inflation and which induces a distinct B-mode polarization pattern in the CMB.\\cite{DurrerBook} Moreover, Inflation typically ends in a violent process at preheating,\\cite{preheating} where large density waves collide at relativistic speeds generating a stochastic background of GW~\\cite{GWpreh} with a non-thermal spectrum characterized by a prominent peak at GHz frequencies for GUT-scale models of inflation (or at mHz-kHz for low scale models of inflation), and an amplitude proportional to the square of the mass scale driving/ending inflation. Such a background could be detected with future GW observatories like Adv-LIGO~\\cite{ligo}, LISA~\\cite{lisa}, BBO~\\cite{bbo}, etc. Furthermore, if inflation ended with a global phase transition, like in certain scenarios of hybrid inflation, then there is also a GWB due to the continuous self-ordering of the Goldstone modes at the scale of the horizon,\\cite{Krauss} which is also scale-invariant on subhorizon scales,\\cite{FFDGB} with an amplitude proportional to the quartic power of the symmetry breaking scale, that could be detectable with laser interferometers as well as indirectly with the B-mode polarization of the CMB.\\cite{GBDFFK}  \\begin{figure} \\hspace{-0.6cm} \\psfig{figure=energyGW.pdf,height=6.5cm} \\hspace{-1.0cm} \\psfig{figure=boundsGW.pdf,height=8.0cm,width=10cm} \\caption{LEFT: The time evolution of the GW energy density during the initial stages of preheating after hybrid inflation, from Ref.[6]. Note the three stages of tachyonic growth, bubble collisions and turbulence. RIGHT: The observational bounds on GW. Flat spectra correspond to GUT Inflation (dashed line) and global phase transitions (continuous lines), while the peaks at the end of the latter spectra correspond to preheating at high scales. \\label{fig:GWpreh}} \\end{figure}   Gravitational waves produced during inflation arise exclusively due to the quasi-exponential expansion of the Universe~\\cite{MukhanovBook}, and are not sourced by the inflaton fluctuations, to first order in perturbation theory. They have an approximately scale invariant and Gaussian spectrum whose amplitude is proportional to the energy density during inflation. GUT scale inflation has good chances to be discovered (or ruled out) by the next generation of CMB anisotropies probes, Planck~\\cite{Planck} and CMBpol~\\cite{CMBpol}, see Fig.~1 for present bounds. At the end of inflation, reheating typically takes place in several stages. There is first a rapid (explosive) conversion of energy from the inflaton condensate to the fields that couple to it. This epoch is known as preheating~\\cite{preheating} and occurs in most models of inflation. It can be particularly violent in the context of hybrid inflation, where the end of inflation is associated with a symmetry breaking scenario, with a huge range of possibilities, from GUT scale physics down to the Electroweak scale. Gravitational waves are copiously produced at preheating from the violent collisions of high density waves moving and colliding at relativistic speeds~\\cite{GWpreh}, see Fig.~1. The GW spectrum is highly peaked at the mass scale corresponding to the symmetry breaking field, which could be very different from the Hubble scale.  \\begin{figure} \\psfig{figure=CBl_defects_inf_800_noise_lense.pdf,height=5.1cm} \\hspace{0.2cm}\\psfig{figure=CBth_defects_inf_lin_400.pdf,height=4.9cm} \\caption{LEFT: The local $\\tilde{B}$-polarization power spectra for tensor perturbations from inflation, cosmic strings, textures and the large-$N$ limit of thenon-linear sigma-model. All spectra are normalized such that they make up 10\\% of the observed temperature anisotropy at $\\ell=10$. Note that the lensed EE modes (red dashed line) contribute as colored noise to the local-BB angular power spectrum. RIGHT: The local $\\tilde{B}$-polarization angular correlation functions for $\\theta<1^o$ for inflation, cosmic strings, textures and the large-$N$ limit of NL sigma model. \\label{fig:CB}} \\end{figure}    In low-scale models of hybrid inflation it is possible to attain a significant GWB at the range of frequencies and sensitivities of LIGO or BBO. The origin of these GW is very different from that of inflation. Here the space-time is essentially static, but there are very large inhomogeneities in the symmetry breaking (Higgs) field due to the random spinodal growth during preheating. Although the transition is not first order, ``bubbles\" form due to the oscillations of the Higgs field around its minimum. The subsequent collisions of the quasi-bubble walls produce a rapid growth of the GW amplitude due to large field gradients, which source the anisotropic stress-tensor~\\cite{GWpreh}. The relevant degrees of freedom are those of the Higgs field, for which there are exact analytical solutions in the spinodal growth stage, which later can be input into lattice numerical simulations in order to follow the highly non-linear and out-of-equilibirum stage of bubble collisions and turbulence. However, the process of GW production at preheating lasts only a short period of time around symmetry breaking. Soon the amplitude of GW saturates during the turbulent stage and then can be directly extrapolated to the present with the usual cosmic redshift scaling. Such a GWB spectrum from preheating would have a characteristic bump, worth searching for with GW observatories based on laser interferometry, although the scales would be too small for leaving any indirect signature in the CMB polarization anisotropies. Moreover, the mechanism generating GW at preheating is also active in models where the SB scenario is a local one, with gauge fields present in the plasma, and possibly related to the production of magnetic field flux tubes~\\cite{PMFpreh}. In such a case, one could try to correlate the GWB amplitude and the magnitude and correlation length of the primordial magnetic field seed.    In the case that inflation ends with a global or local symmetry breaking mechanism, then there generically appear cosmic defects associated with the topology of the vacuum manifold. For instance, global cosmic strings are copiously produced during preheating if the Higgs field is a complex scalar with a U(1) global symmetry~\\cite{preheating}. In principle, all kinds of topological and non-topological defects could form at the end of inflation and during preheating. Such defects will have contributions to all three different metric perturbations: scalar, vector and tensor, with similar amplitudes~\\cite{defrev}. In a recent work~\\cite{FFDGB}, we analyzed the production of GW at preheating for a model with O(N) symmetry. The dynamics at subhorizon scales was identical to that of the usual tachyonic preheating. However, in this model even though the Higgs potential fixed the radial component to its vev, there remained the free (massless) Goldstone modes to orient themselves in an uncorrelated way on scales larger than the causal horizon. In the subsequent evolution of the Universe, as the horizon grows, spatial gradients at the horizon will tend to reorder these Goldstone modes in the field direction of the subhorizon domain. This self-ordering of the fields induces an anisotropic stress-tensor which sources GW production. In the limit of large N components, it is possible to compute exactly the scaling solutions, and thus the amplitude and shape of the GWB spectrum. It turns out that the GWB has a scale-invariant spectrum on subhorizon scales~\\cite{Krauss} and a $k^3$ infrared tail on large scales~\\cite{FFDGB}, which can be used to distinguish between inflation and these non-topological defects.  \\begin{figure} \\psfig{figure=SNR_defects_inf_400_01_log.pdf,height=5.3cm} \\hspace{0.5cm}\\psfig{figure=SNR_defects_inf_400_ratio.pdf,height=5cm} \\caption{LEFT: The signal-to-noise ratio as a function of the normalized polarization sensitivity, for inflation, cosmic strings, textures and the large-$N$ limit of the non-linear sigma-model. RIGHT: The relative signal-to-noise ratio for defect models versus inflation for local (continuous lines) and non-local (dashed lines) B-modes. \\label{fig:SNR}} \\end{figure}   Apart from the IR tail, the main difference between inflationary and global defect contributions to the CMB anisotropies arises from the fact that defects generically contribute with all modes: scalar, vector and tensor modes, with similar amplitudes, while inflationary tensor modes could be negligible if the scale of inflation is well below the GUT scale. Since (curl) B-modes of the polarization anisotropies only get contributions from the vector and tensor modes, the detection of the B-mode from inflation may be challenging, and dedicated experiments like Planck and CMBpol have been designed to look for them. On the other hand, defects' contribution to the temperature anisotropies have a characteristic smooth hump in the angular power spectrum, which allows one to bound their amplitude (and thus the scale of symmetry breaking) below $10^{16}$ GeV.~\\cite{defCMB} However, the contribution to the B-mode coming from defects have both tensor and vector components, and the latter can be up to ten times larger than the former, and actually peaks at a scale somewhat below the horizon at last scattering (in harmonic space the corresponding multipole is $\\ell\\sim1000$).  In a recent paper~\\cite{GBDFFK} we analyzed the possibility of disentangling the different contributions to the B-mode polarization coming from defects versus that from inflation. The main difficulty, for both defects and tensor modes from inflation, is that the B-mode power spectrum is ``contaminated\" by the effect of lensing from the intervening matter distribution on the dominant E-mode contribution on similar angular scales. Using the temperature power spectrum to determine the underlying matter perturbation from evolved large scale structures responsible for CMB lensing, it is possible to engineer an iterative scheme to clean the primordial B-modes from lensed E-modes~\\cite{SeljakHirata}. This procedure leaves a significantly smaller polarization noise background $\\Delta_{P,{\\rm eff}}$ which allows one to detect the GW background at high confidence level (3-$\\sigma$) if the scale of inflation or that of symmetry breaking is high enough. What we realized is that the usual E- and B-modes used for computing the angular power spectra are complicated non-local functions of the Stokes parameters, involving both partial differentiation and inverse laplacian integration. Such a non-local function requires knowledge of the global polarization on scales as large as the horizon, where the B-mode angular correlation function is negligible and thus prone to large systematic errors. In contrast, the so-called ``local\" \\~E- and \\~B-modes~\\cite{DurrerBook,BZ} can be constructed directly from the Stokes parameters and do not involve any non-local inversion. A direct consequence in this change of variables is that the angular power spectrum of local \\~B-modes has a extra factor $n_\\ell = (\\ell+2)!/(\\ell-1)! \\sim \\ell^4$, which boosts the high-$\\ell$ peak in the defects' power spectra. When compared with the angular correlation function of inflation, it gives a significant advantage to the defects' prospects for detection in future CMB experiments, see Fig.~3 and Table~1.    \\begin{table}[h] \\caption{The limiting amplitude for inflation (r=T/S) and various defects ($\\ep=Gv^2$), at 3-$\\sigma$ in the range $\\theta\\in[0,1^o]$, for Planck ($\\Delta_{P,{\\rm eff}}= 11.2\\,\\mu$K$\\cdot$arcmin), CMBpol-like exp. ($\\Delta_{P,{\\rm eff}}=0.7\\,\\mu$K$\\cdot$arcmin) and a dedicated CMB experiment with $\\Delta_{P,{\\rm eff}}=0.01\\,\\mu$K$\\cdot$arcmin. We have assumed $f_{\\rm sky}=0.7$ for all CMB experiments. \\label{tab:3sigma}} \\vspace{0.4cm} \\begin{center} \\begin{tabular}{|c|ccccc|} \\hline $S/N=3$ & Inflation & Strings & Semilocal & Textures & Large-N \\\\ \\hline \\hline Planck & $0.03$ & $1.2\\cdot10^{-7}$ & $1.1\\cdot10^{-7}$ & $1.0\\cdot10^{-7}$ & $1.6\\cdot10^{-7}$  \\\\ \\hline CMBpol & $10^{-4}$ & $7.7\\cdot10^{-9}$ & $6.9\\cdot10^{-9}$ & $6.3\\cdot10^{-9}$ & $1.0\\cdot10^{-8}$  \\\\ \\hline $\\tilde B$ exp & $10^{-7}$ & $1.1\\cdot10^{-10}$ & $1.0\\cdot10^{-10}$ & $0.9\\cdot10^{-10}$ & $1.4\\cdot10^{-10}$  \\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.4807_arXiv.txt": {
        "abstract": "The centre of our Galaxy harbours a 4 million solar mass black hole that is unusually quiet: its present X-ray luminosity  is more than 10 orders of magnitude less than its Eddington luminosity. The observation of iron fluorescence and hard X-ray emission from some of the massive molecular clouds surrounding the Galactic Centre has been interpreted as an echo of a past $\\mathrm{10^{39}\\ erg\\ s^{-1}}$ flare. Alternatively, low-energy cosmic rays propagating inside the clouds might account for the observed emission, through inverse bremsstrahlung of low energy ions or bremsstrahlung emission of low energy electrons. Here we report the observation of a clear decay of the hard X-ray emission from the molecular cloud Sgr B2 during the past 7 years thanks to more than 20 Ms of INTEGRAL exposure.  This confirms the decay previously observed comparing the 6.4 keV line fluxes measured by various X-ray instruments, but without intercalibration effects. The measured decay time is 8.2 $\\pm$ 1.7 years, compatible with the light crossing time of the molecular cloud core . Such a short timescale rules out inverse bremsstrahlung by cosmic-ray ions as the origin of the X ray emission. We also obtained 2-100 keV broadband X-ray spectra by combining INTEGRAL and  XMM-Newton data and compared them with detailed models of X-ray emission due  to irradiation of molecular gas by (i) low-energy cosmic-ray electrons and (ii) hard X-rays. Both models can reproduce the data equally well, but the time variability constraints and the huge cosmic ray electron luminosity required to explain the observed hard X-ray emission strongly favor the scenario in which the diffuse emission of Sgr B2 is scattered and reprocessed radiation emitted in the past by Sgr A*. The spectral index of the illuminating power-law source is found to be $\\Gamma \\sim 2$ and its luminosity $\\mathrm{1.5-5 \\times 10^{39}\\ erg\\ s^{-1}}$, depending on the relative positions of Sgr B2 and Sgr A$^*$.  Using recent parallax measurements that place Sgr B2 in front of Sgr A$^*$, we find that the period of intense activity of Sgr A$^*$ ended between 75 and 155 years ago. ",
        "introduction": "The supermassive black hole of $\\mathrm{4.3 \\times 10^6\\ M_\\odot}$ \\citep{Ghez08,Gillessen09} that sits at the centre of the Galaxy is remarkably quiescent nowadays with an X-ray luminosity of about $\\mathrm{10^{33-34}\\ erg s^{-1}}$, more than 10 orders of magnitude less than its Eddington luminosity. The brightest X-ray flare recorded so far reached a level of a few  $\\mathrm{10^{35}\\ erg s^{-1}}$ \\citep{Porquet08}. Active Galactic Nuclei have a relatively low duty cycle ($\\mathrm{\\sim 10^{-2}}$, \\citealt{Ho08}), so that they spend most of the time  in a low luminosity state. It is therefore not excluded that Sgr A$^*$ has experienced periods of increased activity in the past. Looking for evidence of such a past activity of Sgr A$^*$, we turn to its surroundings that would probably carry the scars of such an event.  The central molecular zone or CMZ \\citep{Morris96}, although relatively small in size with a radius of about 200 pc, contains about 10\\% of the Galaxy's molecular mass. It is  populated by some of the most massive ($\\mathrm{10^4-10^6\\ M_{\\odot}}$) and dense ($\\mathrm{>\\ 10^5\\  cm^{-3}}$) giant molecular complexes in the Galaxy. Sgr B2 is the densest and most massive among them with a mass of a few $\\mathrm{10^6 M_\\odot}$, and density in its 5 pc core reaching $\\mathrm{10^6\\ cm^{-3}}$. It has an envelope that extends to $\\sim$20 pc with an average density of $\\sim\\mathrm{10^{3}\\ cm^{-3}}$  \\citep{Goldsmith90}.  The CMZ is also a source of non-thermal radiation. ART-P on board GRANAT was the first to detect hard X-ray emission from the region \\citep{Sunyaev93}. The ASCA satellite discovered a strong iron K$\\alpha$ fluorescence line at 6.4 keV with a large equivalent width of 1-2 keV from Sgr B2 \\citep{Koyama96}. Chandra has surveyed the central region of our Galaxy and mapped the distribution of Fe line emission from the molecular clouds \\citep{Park04}. Today, thanks to the improved sensitivity of INTEGRAL/IBIS above 20 keV, the emission of Sgr B2 was resolved from the rest of the Galactic Centre (GC) sources \\citep{Revnivtsev04}.  The features seen in Sgr B2 are not unique among clouds in the CMZ. G0.13--0.13 (also konwn as G0.11-0.11) displays strong 6.4 keV line emission \\citep{Yusef02}, as well as a hard X-ray continuum detected at energies above 50 keV \\citep{Belanger06}. Sgr C shows Fe K$\\alpha$ line emission \\citep{Nakajima09}, however its associated hard X-ray continuum cannot be resolved by current instruments because it lies close to a bright X-ray binary, KS 1741-293 \\citep{deCesare07}.  The origin of this non-thermal emission is controversial.  A likely explanation for the continuum and the strong K$\\alpha$ line, is  Compton scattering and K-shell photo-ionization of neutral or low-ionized iron atoms by an intense X-ray radiation field generated in the GC \\citep{Sunyaev93, Koyama96, Sunyaev98}. Here, the clouds would act as X-Ray Reflection Nebulae (XRN).  Since there has been no report of a sufficiently bright X-ray transient (L $\\ge 10^{37}$ erg/s) close to Sgr B2 in the last 20 years, it has been argued that the origin of the emission should lie in the most central regions and have a luminosity of at least $ 10^{39}$ erg/s for several years. The most probable candidate  is Sgr A$^*$. This, however, requires an increase of more than 6 orders of magnitude compared to its current luminosity\\citep{Koyama96}. It has also been suggested that the Sgr A East supernova remnant interacting with the  GC massive molecular clouds could provide the required luminosity\\citep{Fryer06}.   The propagation of energetic charged particles inside the molecular clouds has also been invoked as an explanation of the observed X-rays from the region. For instance, \\citet{Yusef07} suggested that the density of cosmic rays (CR) is larger in the CMZ than in the rest of the Galaxy. Their hypothesis is based on several results: the observation of diffuse low-frequency radio emission in the central regions \\citep{LaRosa05},  the high estimates of the ionization rate compared to the values obtained in the Galactic disk \\citep{Oka05}, and the detection of hard interstellar emission from the molecular clouds at TeV energies \\citep{Aharonian06}.  Two distinct processes can produce hard X-ray continuum  and neutral iron line emissions. First, low-energy cosmic-ray electrons (LECRE), with energies up to a few hundred keV, diffusing into dense neutral matter will produce nonthermal bremsstrahlung and collisional K-shell ionization \\citep{Valinia00}. The presence of strong non-thermal radio filaments in the vicinity of several molecular clouds makes it reasonable to consider the presence of a large flux of cosmic ray electrons in these regions, see e.g. G0.13-0.13 \\citep{Yusef02} or Sgr C \\citep{Yusef07}. Second, subrelativistic ions propagating in the molecular cloud can radiate through inverse bremsstrahlung in the hard X-ray domain and create iron K$\\alpha$ vacancies. Knock-on electrons generated by collisions of primary ions with the gas also contribute significantly to the total emission \\citep{Dogiel09}.  These subrelativistic ions can also be invoked to explain the putative hot plasma pervading the GC. %   \\citet{Revnivtsev04} have argued that the X-ray emission from Sgr B2 is difficult to explain by LECREs, because the large power that cosmic-ray particles would need to deposit in the molecular gas to account for the X-ray flux would be comparable to the bolometric (mostly infrared) luminosity of the complex. They also put into question the large metallicities implied by the large equivalent width of the line produced by collisional K-shell interaction.   To help resolve this ambiguity, Sgr B2 and the other clouds of the CMZ have been the subject of numerous observation campaigns. In particular, a deep Chandra observation revealed that the 6.4 keV line emission is oriented towards the GC \\citep{Murakami01}. It also allowed the contribution from individual sources in the GMC associated with dust condensations (e.g. Sgr B2 North (N) Middle (M) and South(S) to be distinguished \\citealp{Takagi02}).    Recently, thanks to the monitoring of the region by many instruments, time dependent effects have been observed. First, \\citet{Muno07} have detected time varying diffuse emission in the 4-8 keV range a few arcmin away from  Sgr A$^{\\star}$  comparing Chandra images taken before and after 2004. Since this region is emitting the 6.4 keV line, it is likely that the varying diffuse emission can be linked with iron fluorescence, making this cloud  a likely  X-ray reflection nebula.  More recently, comparing recent observations with various X-ray telescopes (namely Suzaku, XMM-Newton, Chandra), \\citet{Inui09} were able to show a temporal decrease in the K$\\alpha$ line emission from Sgr B2, strengthening the reflection interpretation for this cloud as well.  The hard X-ray continuum carries many clues as regards the origin of the emission: its spectral shape and its temporal behaviour in particular. With its sensitivity and with its very long exposures of the Galactic central region, INTEGRAL can therefore provide invaluable insight into the nature of the hard X-ray emission of the CMZ. We present here the lightcurve and spectrum of the Sgr B2 emission taken over 7 years of observations and show that the Sgr B2 flux has decreased by more than 40\\% since 2003. We then study the combined XMM-Newton/INTEGRAL spectrum obtained from observations prior to 2004 which we use to test LECRE and XRN models. Finally, we discuss the implications of our results in relation to the likely nature of the non thermal emission from the GC molecular clouds. ",
        "conclusions": "INTEGRAL has been regularly monitoring the Galactic Centre for 7 years. Using the complete dataset, we confirm the result from \\citet{Revnivtsev04} that IGR J17475-2822 corresponds to hard X-ray ($>$20 keV) emission from the Sgr B2 molecular cloud. In particular, it is in good positional agreement with the very massive and dense core of the cloud. The larger and more diffuse component might be contributing but its flux should be diluted because of its extent relative to the PSF size. Since the mean density is of the order of $\\mathrm{1-3 \\times 10^3 cm^{-3}}$, the envelope column density is smaller by a factor of a few compared to the core column density (which is close to $\\mathrm{10^{24} cm^{-2}}$), so that the diffuse hard X-ray contribution has to be much smaller than that of the core.  The regular monitoring  shows that the Sgr B2 20-60 keV emission is fading, with a flux decrease of about 40\\% after 7 years of observations. A test of hypothesis shows that a linear decay is favored compared to a constant flux at the 4.8 $\\sigma$ confidence level.  The linear decay characteristic timescale is found to be $\\sim$ 8 yrs,  comparable with the light-travel time across the cloud core.  This result is obtained thanks to the very deep exposure obtained by INTEGRAL on the GC regions and does not suffer from instrument to instrument  intercalibration errors. It is consistent with the results obtained comparing the 6.4 keV fluxes measured by Chandra, XMM-Newton and Suzaku from 2000 to 2005 \\citep{Inui09}.  Cosmic ray models of the iron line and continuum emission are put to a test by this fast variability. Inverse bremsstrahlung emission from subrelativistic ions is ruled out because of the long cooling timescale of these particles in the dense environment of the cloud core. Low energy cosmic ray electrons are not formally ruled out: the lifetime of $\\sim$ 100 keV electrons in the dense matter of Sgr B2 is smaller  than the measured decay time, so that if the cosmic-ray injection had ceased with a similar or smaller timescale, we might obtain a similar flux decrease.  Nevertheless, combining the measured decay with the large electron injection power required (which we fitted to be of the order of $10^{40}$ erg/s) we can exclude this scenario as well.  Sgr B2 is therefore most likely a reflection nebula, as was initially proposed by \\citet{Sunyaev93} and \\citet{Koyama96}.  The location of the illuminating source is yet to be determined unequivocally. As was already noted by several works, an internal or a nearby illuminating source should have been detected in the last decade or so. The most likely explanation is therefore a period of high state of Sgr A$^{\\star}$. The observation of variability of the Fe K$\\alpha$ line in other massive molecular clouds of the CMZ strongly supports this hypothesis \\citep{Ponti10}, since several distinct bright illuminating sources active for at least several years would be required to light all these clouds. They could be illuminated by the same outburst from Sgr A$^*$.  Using a model based on the approach of \\citet{Murakami00}, we tested this scenario and found that the illuminating spectrum must be hard (photon index $\\sim$ 2), which is well in the range of measured Low Luminosity Active Galactic Nuclei (LLAGN) spectral indices \\citep{Gu09,Pellegrini00}. Since parallax measurements place Sgr B2 130 $\\pm$ 60 pc in front of Sgr A$^{\\star}$  \\citep{Reid09}, our measurement implies a mean luminosity of 2-5 $\\times 10^{39}$ erg/s, an increase of more than 5 orders of magnitude compared to the current one placing it among faint LLAGNs \\citep{Ho08,Gu09}. At the end of the outburst the variability of Sgr A$^*$ is constrained to be faster than the measured decay rate, which is comparable to the long term variability observed in M81$^*$ \\citep{Iyomoto01,LaParola04}. The end of the episode of high luminosity in Sgr A*, which is currently being mirrored by Sgr B2, occurred about 100 yrs ago, with the intervening period seeing the output of the central source fall by some 5 orders of magnitude. Assuming that Sgr B2 is at the same distance that Sgr A$^{\\star}$, \\citet{Inui09} have proposed that the luminosity of the latter has been continuously decreasing with a constant decay time for 300 years. The new position implies that Sgr A$^{\\star}$ luminosity had to decrease on a much shorter timescale. The origin of such large swings in the luminosity of the supermassive black hole at the centre of our Galaxy is presently unknown but the subject of intensive on-going research."
    },
    "1005/1005.4036_arXiv.txt": {
        "abstract": "In this work, we investigate the nature of the  host galaxies  of long Gamma-Ray bursts (LGRBs) using a galaxy catalogue constructed from the {\\it Millennium Simulation}. We developed a LGRB synthetic model based on the hypothesis that these events originate at the end of the life of massive stars following the collapsar model, with the possibility of including a constraint on the metallicity of the progenitor star. A complete observability pipeline was designed to calculate a probability estimation for a galaxy to be observationally identified as a host for LGRBs detected by present observational facililties. This new tool  allows us to  build an observable host galaxy catalogue which is required to reproduce the current stellar mass distribution of observed hosts. This observability pipeline predicts that the minimum mass for the progenitor stars should be $ \\sim 75 $ M$_\\odot$ in order to be able to reproduce BATSE observations. Systems in our  observable catalogue are able to reproduce the observed properties of host galaxies, namely stellar masses, colours, luminosity, star formation activity and metallicities as a function of redshift. At $ z >2$, our model predicts that the observable host galaxies would be very similar to the global galaxy population. We found that $\\sim 88$ per cent of the observable host galaxies with mean gas metallicity lower than $0.6\\, {\\rm Z_\\odot}$ have stellar masses in the range  $  10^{8.5}$--$10^{10.3}$M$_\\odot$ in excellent agreement with observations. Interestingly, in our model observable host galaxies remain mainly within this mass range regardless of redshift, since lower stellar mass systems would have a low probability of being observed while more massive ones would be too metal-rich. Observable host galaxies are predicted to preferentially inhabit dark matter haloes  in the range  $10^{11}$--$ 10^{11.5}$M$_\\odot$, with a weak dependence on  redshift. They are also found to preferentially map different density environments at different stages of evolution of the Universe. At high redshifts, the observable host galaxies are predicted to be located in similar environments as the global galaxy population but to have a slightly higher probability to have a close companion. ",
        "introduction": "\\label{int}  Gamma-ray bursts (GRBs) are brief pulses of $\\gamma$-ray radiation observed on average once a day at random directions in the sky. They are the brightest sources in this region of the electromagnetic spectrum, and they have been systematically studied in the past two decades \\citep[e.g.][and references therein]{Mes06}. The origin of GRBs is cosmological, as determined from the measurement of their redshifts \\citep*[e.g.][]{Met97,Blo01,Blo03}, and in many cases, their host galaxies have been identified \\citep*[see][and references therein]{LeF03,Sav09}. Their distances imply the release of large amounts of energy ($\\sim10^{51}\\,{\\rm ergs}$) in a short time-scale, which suggests that the origin of these phenomena could be associated to the accretion of matter onto a compact object \\citep{Woo93,Fry99,McF99,Pan01,Fra01}.  Two populations of GRBs are apparent from the distribution of their duration \\citep{Kou93}. Those lasting less than $2\\,{\\rmn s}$ are known as short GRBs, while longer events are called long GRBs (LGRBs). These are more frequently observed and more precisely located, which makes their properties better known. LGRBs are always found in host galaxies with ongoing star formation activity \\citep*{LeF03,Chr04,Pro04}, and some of them were observed to be associated with Type Ib/c core-collapse supernovae \\citep{Gal98,Hjo03}. These observations suggest that the progenitors of LGRBs are massive stars. Stellar evolution models have tried to provide a consistent scenario where a LGRB can develop. The result is the so called {\\em collapsar model} \\citep*{Woo93,Fry99}, in which this phenomenon is produced during the collapse of a massive star, due to the accretion of part of the envelope onto the recently formed black hole. Nevertheless, the properties of LGRB progenitors (mass, metallicity, rotation velocity, binarity, etc.) are still a matter of discussion. LGRBs have been associated to single massive stars \\citep{Woo93}, perhaps in metallicity biased environments \\citep*{Hir05}, and also to stars in binary systems \\citep{Bis07}.  The importance of understanding the nature of the LGRB progenitors is beyond the interest  of  only  stellar evolution and black hole formation. Given their connection to massive stars and their large luminosities, LGRBs might be a powerful tool to investigate the star formation in the early Universe, at redshifts for which standard tools become ineffective \\citep[e.g.][]{Wij98}. The knowledge of the properties of the stellar progenitors would allow the assessment of possible biases originated when LGRBs are used as tracers of star formation. Several authors suggest that the cosmic LGRB rate does not follow the star formation rate \\citep*{Dai06,Sal07}. The interpretation of their data requires the assumption of a differential evolution of the comoving LGRB rate density with respect to the comoving star formation rate (SFR) density. The same conclusion arises from the work of \\citet{Wol07}, who show that the luminosity function of LGRB hosts differs from that of core-collapse supernovae hosts, which are considered unbiased tracers of star formation. According to these authors, the LGRB host luminosity function can be reproduced by requiring LGRB progenitors to have metal abundances lower than that of the Sun.  An indirect procedure to investigate LGRB progenitors is to characterise the stellar populations of their host galaxies. Given the connection between LGRBs and star formation, galaxy formation models can be used to assess the validity of LGRB progenitor models, comparing their predictions about the properties of these stellar populations to host galaxies observations. Recent works such as that of \\citet{Kau04}, suggest that there is a strong relation between star formation and nuclear activity in a galaxy with its environment, at constant stellar mass. Hence, the environment of host galaxies might constitute an independent probe for progenitor models. Observational works on host galaxies environment are still inconclusive on whether host galaxies inhabit regions with certain characteristics \\citep{Fyn02,Jak05,Tho08}. \\citet{Bor04} suggest that host galaxies are field galaxies, while \\citet*{Wai07} find that a considerable fraction of their sample of 42 host galaxies shows evidence of interaction with other galaxies. The question of whether there is a connection between the occurrence of a LGRB and the local density of galaxies remains unanswered.  Since galaxy formation is a highly non-linear process, the properties of host galaxies predicted by different progenitor scenarios are better studied by means of cosmological numerical simulations \\citep{Kat91,Nav93,Mos01,Spr03,Sca05,Sca06}. \\citet*{Cou04} use hydrodynamical simulations of structure formation to identify a galaxy population whose properties reproduce the observed ones. According to these authors, the luminosity distribution of host galaxies is reproduced if galaxies are required to have high star formation efficiency. However, their simulations do not make any prediction on the properties of the LGRB progenitors. \\citet{Nuz07} developed a Monte Carlo code to simulate the production of LGRBs in hydrodynamical simulations of galaxy formation, assuming that their progenitors are massive stars as proposed by the collapsar model. These authors follow the  LGRB production  as the structure forms and evolves in the Universe. Their results suggest that LGRB progenitors would be low metallicity stars ($Z < 0.3 Z_\\odot$), and hence LGRBs would be biased tracers of star formation principally at low redshift. However, their simulations explore a small volume of $(10 {h^{-1}\\,\\rm Mpc})^3$  dominated by field galaxies. Recently, \\citet{Cam09}  have constructed a simulated population of host galaxies based on the semi-analytic model of galaxy formation of \\citet{DeL07} applied to cosmological simulations developed by \\citet{Wan08}. These simulations describe a larger volume of the Universe (a square box of $125 {h^{-1}\\,\\rm Mpc}$ in a side). \\citet{Cam09} explore three LGRB progenitor models based on the collapsar model, one of them assuming a given stellar mass value and an age cut-off and  the other two also  including  metallicity cut-offs. These authors find that models with very low metallicity progenitors ($Z < 0.1 Z_\\odot$) could explain the luminosities, colours and metallicities of the observed host galaxies. Their results support previous claims of LGRBs being biased tracers of the star formation.  Clearly, the use of larger simulations with better resolution and more detailed descriptions of the physical processes driving the dynamical, chemical and star formation histories of galaxies allows a better modelling of the host galaxies of LGRBs.  However, an important aspect that should not be disregarded is that, for a proper comparison with observed host galaxies, the effects of the detectability of these galaxies must be taken into account. And this is the main contribution of our work where we developed an observational pipeline which allows us to mimic, at least in part, biases affecting the observations. With this new tool we build an observable host galaxy catalogue which can be more fairly compared to current available observations.  In this work, we develop a semi-analytical model for the  host galaxies of LGRBs adopting the collapsar model for the  progenitor stars. We apply it to one of the largest cosmological simulations available, the {\\em Millennium Simulation} \\citet{Spr05a}. The {\\em Millennium Simulation} follows the evolution of dark matter in a box of $500 {h^{-1}\\,\\rm Mpc}$ in a side, hence providing a cosmologically representative volume a factor of 64 larger than that of \\citet{Wan08} and factor of $\\sim 10^5$ larger than that of \\citet{Nuz07}. We take the galaxy catalogue built up by \\citet{DeL07}, who describe the evolution of baryonic matter including star formation, active galactic nuclei and supernova feedback. Using the star formation and chemical properties of galaxies in this catalogue to implement different LGRB progenitor scenarios, we determine the properties of the overall host galaxy population and its environment. We compute the detectability of the host galaxies in order to compare the predictions of these scenarios to the observations compiled by \\citet{Sav09}.  This work is organized as following. In Section~\\ref{sim}, we briefly describe the cosmological simulation and semi-analytic galaxy catalogue used in our work. In Section~\\ref{mod}, we develop our implementation of LGRB progenitor scenarios. Sections~\\ref{hgs} and \\ref{env} show our results on host galaxy properties and environment, respectively. Finally, Section~\\ref{con} presents our conclusions.   \\section[]{Millennium Simulation} \\label{sim}  The {\\em Millennium Simulation} \\citep{Spr05a} is one of the largest cosmological simulations of structure formation publicly available. It describes the evolution of the cold dark matter while the baryonic matter can be included via semi-analytic models on top of the numerical simulation. The {\\em Millennium Simulation} describes the formation of structure tracking $\\sim10^{10}$ dark matter particles of mass $8.6 \\times 10^8 h^{-1}\\,{\\rm M_{\\odot}}$ distributed in a cubic region of 500 $h^{-1}\\,{\\rm Mpc}$ side, using the $\\Lambda{\\rm CDM}$ cosmogony. The adopted cosmological parameters are $\\Omega_{\\rm m}=0.25$, $\\Omega_{\\rm b}=0.045$, $\\Omega_{\\Lambda}=0.75$, $\\sigma_8=0.9$ and $n=1$, where the Hubble constant is $H_0=100{h\\,\\rm km} \\,{\\rm s}^{-1}\\,{\\rm Mpc}^{-1}$ with $h=0.73$. These parameters are chosen in consistency with results obtained by the joint analysis of the Two-degree Field Galaxy Redshift Survey (2dFGRS\\footnote{http://www.mso.anu.edu.au/2dFGRS/}) and the Wilkinson Microwave Anisotropy Probe (WMAP) data \\citep{San06}.  The simulation was performed using a modified version of GADGET-2 hydrodynamical code \\citep{Spr05b}. A {\\em friends-of-friends} (FOF) algorithm identified non-linear dark matter haloes within the simulation in an automatic manner. Two particles belong to the same halo if their separation is less than $0.2$ times the mean separation between particles. Only groups of more than $20$ particles are identified as haloes. This restriction sets a lower limit on the mass of a halo of $1.72 \\times 10^{10} h^{-1}\\,{\\rm M_{\\odot}}$. Substructures orbiting within each virialized halo are identified applying a SUBFIND algorithm \\citep{Spr01}. Each dark matter halo has a central (Type 0) galaxy and one or more satellite galaxies. Satellites were at some point central galaxies of smaller haloes which suffered a merger with the halo their currently inhabit. There are two types of satellite galaxies, Type 1 galaxies are located at the centre of a subhalo associated with a FOF group, while Type 2 galaxies have lost their dark matter subhalo after falling  onto a more massive halo.  Mean properties of synthetic galaxies (e.g. SFR, stellar mass, gas mass, metallicity) are obtained by applying semi-analytic models to the structure formation simulation \\citep{Whi91,Col91,Lac91}. Essentially, semi-analytical models describe the collapse of baryonic matter following dark matter haloes. After collapsing, the gas cools and gravitational instabilities produce episodes of star formation. To reproduce observations, models incorporate photoionization processes in the intergalactic medium, the growth of supermassive black holes during galactic mergers, supernova and AGN feedback, star formation rate enhancement in mergers, and formation of heavy elements for each formed stellar population. The galaxy catalogue we use in this work is that of \\citet{DeL07}.  There is recent evidence for an excessive reddening of Type 2 galaxies \\citep*{Wei06,Per09} which might be due to the poor physical treatment of this type of galaxies. In order to avoid spurious trends, we exclude Type 2 galaxies from our analysis in Section~\\ref{mod} and Section~\\ref{hgs}. However, Type 2 galaxies ought to be included in the environmental analysis performed in Section~\\ref{env} in order to correctly trace the underlying mass distribution.  Hence, the catalogue of \\citet{DeL07} provides us with the spatial galaxy distribution, their stellar masses, dark matter haloes, star formation activity, colours, luminosities and mean metallicities. Regarding the latter, the public catalogue makes available the mean metallicity of the cold gas component and of the stellar population as a whole (i.e. averaged over new and old stars in a galaxy). Therefore, as a proxy for the metallicity of the LGRB progenitors, we take the mean metallicity of the cold gas component at the time these stars were born. Because the \\citet{DeL07} model assumes instantaneous recycling, this metallicity is slightly higher than that of the newly born stars, but this correction is negligible compared to the uncertainties in the measured metallicities  \\citep[e.g.][]{Cam09}.   \\section[]{Scenarios for LGRB progenitors} \\label{mod}   \\subsection{Intrinsic LGRB rate}  We consider two scenarios for LGRB production, both of them based on the collapsar model \\citep{Woo93,McF99,Fry99}. In our scenario I we take as LGRB progenitors all stars above a certain minimum mass $m_{\\rm min}$, with no other restriction whatsoever. In this scenario LGRBs are unbiased tracers of star formation. Accordingly, we obtain the intrinsic LGRB rate in a given galaxy $g$ at a particular redshift $z$ as  \\begin{equation} \\label{rateint} r_{\\rm GRB}(g,z)= {\\rm SFR}(g,z) \\frac {\\int_{m_{\\rm min}}^{100\\,{\\rm M_{\\odot}}} \\xi(m)\\,dm}{\\int_{0.1\\,{\\rm M_{\\odot}}}^{100\\,{\\rm M_{\\odot}}} m \\xi(m)\\,dm}, \\end{equation}  \\noindent where $0.1\\,{\\rm M_{\\odot}}$ and $100\\,{\\rm M_{\\odot}}$ are the lower and upper mass cut-offs of the IMF $\\xi(m)$ given by \\citet{Cha03}. According to Eqn.~\\ref{rateint}, the production of LGRBs is not delayed with respect to the starburst that created the progenitor stars, which is justified because in the collapsar model only massive stars are LGRB progenitors.  In our scenario II, we assume that only massive stars ($m > m_{\\rm min}$) below some metallicity threshold ($Z_{\\rm C}$) are able to produce LGRBs, as in the progenitor models of \\citet{Hir05} or \\citet*{Yoo06}.  As explained before, we take the mean metallicity of cold gas in each galaxy as representative of the metallicity of the progenitor stars. The consequences of this hypothesis will be discussed in later sections. In this scenario, the intrinsic LGRB rate is given by Eqn.~\\ref{rateint} for galaxies $g$ with metallicity $Z_g < Z_{\\rm C}$, and is null for those with $Z_g \\geq Z_{\\rm C}$. Three realizations of this scenario were computed, adopting $Z_{\\rm C} = 0.1,\\, 0.3,\\, 0.6\\, {\\rm Z_{\\odot}}$ (scenarios II.1, II.2 and II.3, respectively).   \\subsection{Observed LGRB rate} \\label{obsrate}  Given a metallicity threshold, we adjusted the only free parameter of our scenarios ($m_{\\rm min}$) to reproduce the LGRB rate measured by the BATSE experiment. For this purpose, we took the off-line GRB search of \\citet{Ste01}, which detected 3475 GRBs with $T_{90} > 2\\,{\\rm s}$ during a live-time of $6.37\\,{\\rm yr}$ (70 per cent of the $9.1\\,{\\rm yr}$ that BATSE was active), scanning 67 per cent of the sky. This translates into a full-sky LGRB rate of $814\\,{\\rm yr}^{-1}$ above the threshold of the off-line search. The choice of this particular experiment was motivated by its good statistics and the availability of an accurate model for its detection efficiency.  To obtain the observable rate predicted by a given scenario, we first compute the comoving LGRB rate density for each redshift as  \\begin{equation} \\label{e:comovrate} n_{\\rm GRB}(z) = V^{-1} \\sum_g r_{\\rm GRB}(g,z), \\end{equation}  \\noindent where $V = (500 \\ h^{-1}\\ {\\rm Mpc})^3$ is the comoving volume of the {\\it Millennium Simulation}. The full-sky observed LGRB rate is then  \\begin{equation} \\label{ratebatse} R_{\\rm GRB} = \\int_0^{z_{\\rm max}}\\frac{n_{\\rm GRB}(z)}{(1+z)} p_{\\rm det,BATSE}(z) \\frac{dV}{dz}, \\end{equation}  \\noindent where  \\begin{equation} \\label{dvdz} \\frac{dV}{dz} = \\frac{4 \\pi c d_L^2(z)}{H(z)(1+z)^2} \\end{equation}  \\noindent is the derivative of the comoving volume with respect to $z$ at fixed solid angle, $z_{\\rm max}$ is the maximum redshift of the simulation, $d_L$ the luminosity distance and $H$ the Hubble constant at redshift $z$, and $p_{\\rm det,BATSE}$ the probability of detecting a LGRB at redshift $z$ with BATSE (i.e. the probability that a particular LGRB has a peak flux above the experiment threshold). The integration was performed numerically, and the described procedure iterated over  $m_{\\rm min}$, until agreement with the observed rate was attained.  The value of $p_{\\rm det,BATSE}(z)$ depends on the LGRB luminosity function and spectrum. It was computed using a Monte-Carlo scheme to simulate, at a given redshift, a large population of LGRBs with different luminosities and spectra. Then, their photon fluxes in the BATSE energy band were estimated. For each LGRB in this population, a second Monte-Carlo procedure rejected those events which would be undetectable, taking into account the trigger efficiency of the \\citet{Ste01} off-line search. If the LGRB emission were isotropic, the fraction of retained LGRBs would directly give the detection probability due to the off-line search threshold. To account for the beamed emission of LGRBs \\citep[e.g.][]{Yon05}, and assuming that the distribution of jet opening angles is independent of the LGRB luminosity, spectrum and redshift, we obtain $p_{\\rm det,BATSE}$ by multiplying this fraction by the mean beaming fraction of the jets $p_{\\rm jet} = 10^{-3}$. The luminosity function and spectral parameters distributions were taken from \\citet{Dai06}.  This procedure not only allows us to mimic the observational process but also has provided us with estimates of the minimum mass for the progenitor stars. The values of $m_{\\rm min}$ obtained, in solar masses, are $91.4 \\pm 0.1$, $13.9 \\pm 0.2$, $44.3 \\pm 0.4$, and $76.0 \\pm 0.3$ for scenarios I, II.1, II.2 and II.3, respectively; the uncertainties reflect Poissonian errors in the number of observed LGRBs (see also Table 1).  \\begin{table} \\begin{center} \\caption{Main characteristics of our scenarios for the simulated host galaxies of LGRBs. Column (1) gives the name of the scenario.  Column (2) lists the minimum stellar mass obtained for progenitor stars of LGRBs to reproduce BATSE observations, adopting \\citet{Cha03} IMF. Column (3) gives the cold gas maximum metallicity cut-off. Column (4) shows ratio between  the percentage of simulated host galaxies over the mass  range $\\approx 10^{8.5-10.3} M_\\odot$  and the corresponding value obtained from the sample of \\citet{Sav09}.} \\label{tab1} \\begin{tabular}{|c|c|c|c|}\\hline {Scenarios}  & { $m_{\\rm min}$ (${\\rm M_\\odot}$)} & { $Z_{\\rm C}$}  &  {R$_{\\rm M}$}  \\\\\\hline I   &  $91.4 \\pm 0.1$   & - & 0.82   \\\\ I.1 & $13.9 \\pm 0.2$    & 0.1& 0.02\\\\ II.2 & $44.3 \\pm 0.4$ &  0.3  & 0.51\\\\ II.3 & $76.0 \\pm 0.3$   & 0.6 &  1.03\\\\\\hline \\end{tabular} \\end{center} \\vspace{1mm} \\end{table}    \\subsection{Host galaxies}  A meaningful definition of a LGRB host galaxy in our scenarios is not as straightforward as it might seem at first sight. The na\\\"ive definition of a host galaxy, at a given redshift $z$, as being any galaxy $g$ with a LGBR rate of $r_{\\rm GRB}(g,z) > 0$ (or identically, ${\\rm SFR}(g,z) > 0$) is not useful because the resulting host population would include galaxies with arbitrarily low LGRB rates. Low rates are better understood in statistical terms, as very low probabilities of producing a LGRB per unit time. This implies that the corresponding galaxies have a low probability of being observed as host galaxies. Hence, this na\\\"ive definition would generate a host galaxies sample biased to low SFR galaxies. To fix this problem, the definition of a host galaxy could be based on the number of LGRBs $N(g,z) = r_{\\rm GRB}(g,z) \\Delta t$ produced in each galaxy during a time interval $\\Delta t$ in its rest frame, so that host galaxies are only those with $N(g,z) \\geq 1$, as in \\citet{Cam09}. However this cut-off, and the resulting population, would be dependent on the more or less arbitrary choice of $\\Delta t$.  Given the above arguments, we preferred instead a probabilistic approach, defining the likelihood of a galaxy being observed as a LGRB host. This likelihood is then used as a weight to compute the properties of the observable host galaxies sample, which would be in this way comparable to the observed sample. For a galaxy to be detected as a host at least one LGRB must be detected within it by a high-energy observatory. Then, the galaxy itself must be detected, usually by optical/NIR telescopes. We would treat the biases introduced by each of these observations separately.  To model the first bias, we compute the probability of detecting at least one LGRB within it by any of the high-energy observatories monitoring GRBs. Given that event detection processes follow Poissonian statistics, if $R_{{\\rm GRB},i}(g,z)$ is the contribution of galaxy $g$ at redshift $z$ to the LGRB rate observed by experiment $i$, the probability of observing at least one LGRB in a galaxy is  \\begin{equation} \\label{phg} p_{\\rm HG}(g,z) = 1 - \\exp \\left(-\\sum_i R_{{\\rm GRB},i} \\Delta T_i \\right), \\end{equation}  \\noindent where $\\Delta T_i$ is the time interval of the LGRB search conducted by experiment $i$. Eqn.~\\ref{phg} assumes that the searches by different observatories are independent. Previous works estimate the intrinsic LGRB rate in a typical galaxy to be of the order of $10^{-3}\\,{\\rm yr}^{-1}$ \\citep{Fry99}, hence the LGRB rate observed in the galaxy by any mission must be lower than this value. Given that current searches for LGRBs have spanned a few years, $\\sum_i R_{{\\rm GRB},i} \\Delta T_i \\ll 1$, and $p_{\\rm HG}(g,z) \\simeq \\sum_i R_{{\\rm GRB},i} \\Delta T_i$, which means that the likelihood of a galaxy being observed as a host is proportional to its contribution to the mean {\\em observed} number of LGRBs,   \\begin{equation} \\label{rateobsgal} p_{\\rm HG}(g,z) \\simeq  \\frac{r_{\\rm GRB}(g,z)}{V (1+z)} \\frac{dV}{dz} \\sum_i p_{{\\rm det},i}(z) \\Delta T_i \\frac{\\Delta \\Omega_i}{4 \\pi}, \\end{equation}   \\noindent where $\\Delta \\Omega_i$ is the sky coverage of the experiment $i$. Note that, for a fixed $z$, cosmological and observatory dependent factors in Eqn.~\\ref{rateobsgal} become constant, independently of the number of observatories considered, and cancel out in the normalization of the weights. This means that when computing mean HG properties such as mass or SFR as a function of $z$, the first bias can be modeled simply by weighting the corresponding properties of each galaxy by its {\\em intrinsic} LGRB rate.  In the computation of the integrated properties of the whole population of host galaxies (like the integrated mass distribution, for example), also the effects of volume variation, time dilation and detectability of the different observatories must be taken into account (Pellizza et al. in preparation). In this case it is better to compute first the mass distribution observed by each experiment at each redshift $z$, weighting the galaxies in the corresponding snapshot by $p_{\\rm HG}(g,z)$. Second, the integration in $z$ can be performed for each experiment, and the resulting distribution can be normalized. This has the advantage of avoiding the use of the values of $\\Delta T_i$ and $\\Delta \\Omega_i$, which are poorly known for some of the observatories that detected the LGRBs in the sample of \\cite{Sav09}. Only $p_{{\\rm det},i}(z)$ for each experiment is needed, which is computed as described in Sect.~\\ref{obsrate} for $p_{\\rm det,BATSE}(z)$. Finally, these distributions can be combined into a single one by adding them, previously scaled to the number of LGRBs detected by each observatory. The relevant data for computing the detectability were taken from \\citet{Ste01} for BATSE, \\citet{Gue07} for {\\it Swift}, \\citet{Fro09} for {\\it Beppo-SAX}, \\citet{Pel08} for {\\it HETE-2} and \\citet{Hur92} for {\\it Ulysses}. These experiments detected 36 of the 38 LGRBs in the sample of \\citet{Sav09}. For statistical reasons, only experiments that detected at least five LGRBs in the sample were considered. Konus-{\\it Wind} and {\\it NEAR}-XGRS were discarded because all their LGRBs were observed also by other experiments above. To be consistent, the observed mass distribution to which the model predictions were compared was constructed from the data of \\citet{Sav09} by adding the individual distributions observed by each experiment.  The second bias is more difficult to model. Given that the search for host galaxies is guided by the discovery of the LGRBs themselves and usually done with a variety of different telescopes and detectors, in different bands of the electromagnetic spectrum and with different sensitivities, the biases introduced are unclear. Galaxies with surface brightness below the detectability threshold of available instruments could introduce a bias towards low metallicity galaxies. An extra problem could be caused by dust obscuration which could affect the detection of the  afterglows and produce a bias towards high metallicity galaxies \\citep{Fyn09}. Hence, at least these two effects might combine themselves to determine the detectability of a host galaxy. Considering also the fact that we have only  access to the public galaxy catalogue of the {\\it Millennium Simulation} which provides only mean global properties and magnitudes, we adopt the observed integrated stellar mass distribution as a tool to apply the combined effect of observational biases to the simulated galaxy sample. We use the fact that  the observed integrated stellar mass distribution has been affected by observational biases although we cannot disentangle their individual effects. Hence, we require the simulated LGRB hosts to reproduce the observed stellar mass distribution in order to build up the observable simulated LGRB hosts. The procedure is explained in detail in next section. Note that, hereafterin, we will discuss the trends of the observable simulated hosts and the general galaxy population. The former can be compared to the current observed hosts, but, if this later sample changes due to better or different observational techniques, our simulated sample should be also consistently modified to match the new observed stellar mass distribution.  \\section[]{LGRB host properties} \\label{hgs}  The largest and most comprehensively uniform sample of host galaxy properties available at present is that compiled by \\citet{Sav09}\\footnote{Data available at http://www.grbhosts.info}. Stellar masses, star formation rates, metallicities, absolute magnitudes and colours of 46 observed host galaxies up to $z \\sim 3$ were obtained by these authors comparing their spectral energy distributions to those of synthetic stellar populations, using the method described in \\citet{Gla04}.  A key point of our models is that we require the stellar mass distribution of the predicted host galaxies to  match that of the observed host galaxies.  The stellar mass is adopted as the property to be reproduced by the models since it is now widely accepted that stellar mass is a more fundamental quantity for galaxies than luminosity \\citep[e.g.][]{Kau04}. It could be possible that those LGRB events with no detected hosts occured in very low surface brightness galaxies and hence, with low stellar masses. As mentioned before the effects of dust could also prevent the detection of events in high metallicity (and dusty) galaxies. In our models these  host galaxies exist (and are part of the global galaxy population) but the observability cut-off has been determined by using the  current observed stellar mass distribution. Fig.~\\ref{mass} shows the distribution of stellar masses of observed host galaxies constructed from the data of \\citet{Sav09}, together with those predicted by our scenarios.  As it can be seen from Fig.~\\ref{mass}, the scenario which best reproduces the observed stellar mass distribution is that with $Z_{\\rm C}=0.6\\, {\\rm Z_{\\odot}}$ (scenario II.3, see also Table 1). Lower metallicity thresholds predict lower observed stellar masses for the host galaxies, while no metallicity threshold (scenario I) predicts larger ones. In the sample of \\citet{Sav09}, we find that $85$ per cent of the studied galaxies have stellar masses over the range $\\approx 10^{8.5-10.3} M_\\odot$ \\citep[see also][]{Cas08}. The probability of getting a host galaxy within this stellar mass range in scenario II.3 is 88 per cent, while for scenarios I, II.1 and II.2 it is 70, 2, and 43 per cent,  respectively. Then we conclude that scenario II.3 predictions agree fairly well with observations while others fail, and therefore in the following sections,  we will focus only on this scenario.  We stress the fact that its predictions include the effects of host galaxies observability, hence, a proper comparison with observations can be made. In order to contribute to the understanding of the nature of LGRB host galaxies, we will also compare these predictions with the properties (not weighted by host galaxy observability) of both the sample of all galaxies with mean cold gas metallicities below $0.6\\, {\\rm Z_{\\odot}}$ (hereafter low metallicity sample) and the complete galaxy population of the catalogue of \\citet{DeL07}.  \\subsection{Stellar mass}     \\begin{figure} \\includegraphics[width=0.45\\textwidth]{fig1.ps} \\caption{Stellar mass distribution of host galaxies in the sample of \\citet{Sav09} (shaded in gray), together with those predicted by scenarios I (dashed-triple dotted line), II.1 (dotted line), II.2 (dashed-dotted line) and II.3 (solid line).} \\label{mass} \\end{figure}  \\begin{figure} \\includegraphics[width=0.45\\textwidth]{fig2.ps} \\caption{Mean  stellar mass of the  host galaxies as a function of redshift predicted by scenario II.3 (dotted line) with its standard deviation  (shaded grey band). Open diamonds represent the  observed host galaxies of \\citet{Sav09}, while filled circles correspond to their mean values in redshift intervals of 0.5. For comparison, we also include the corresponding mean trends for the low metallicty sample (dashed line) and the complete galaxy population (solid line).} \\label{massred} \\end{figure}  As a first step towards understanding the nature of host galaxies, we analyse their stellar masses as a function of redshift. As shown in Fig.~\\ref{massred}, the mean stellar mass of the host galaxies as a function of redshift predicted by scenario II.3 reproduces the observed mean trend quite well. From this figure, we can also see that host galaxies are, on average, more massive than galaxies in the complete galaxy population while the latter are more massive than those in the low metallicity sample. This can be understood taking into account that the host galaxy observability is a strong function of the star formation rate, and that the complete galaxy population in the catalogue of \\citet{DeL07} follows a mass-metallicity relationship \\citep[e.g.][]{DeR09}. Then, the cut-off adopted for the  mean cold gas metallicity to reproduce the observed stellar mass distribution implies a cut-off in stellar mass since low metallicity galaxies are, on average, less massive than the general galaxy population (see also Fig. \\ref{mass}). However, as the  observability of a host galaxy depends strongly on its star formation activity and most of the small galaxies have low star formation rates, the observable host galaxies tend to be, on average, the  more massive ones among them. As a result, our observable sample tends to be populated by systems more massive than those in  the low metallicity sample or in  the complete galaxy population.  \\subsection{Star formation rate}  In Fig.~\\ref{gsfr}, we display the mean SFR of  host galaxies as a function of redshift predicted by scenario II.3, together with the corresponding mean values for the low metallicity sample and the complete galaxy population. As it can be seen, the prediction of scenario II.3 reproduces very well the behaviour of the observed host galaxies. These have higher SFR than the mean of the complete galaxy population, and much higher than that of the low metallicity sample. The good agreement between our scenario II.3 and observations suggests that the observed host galaxies are biased towards galaxies with stellar masses in the range $10^{9-10} {\\rm M_\\odot}$, high star formation activity and relatively low gas metal content, compared to the mean properties of the complete galaxy population at a given $z$.  \\begin{figure} \\includegraphics[width=0.45\\textwidth]{fig3.ps} \\caption{Mean  SFR of the host galaxies as a function of redshift predicted by scenario II.3 (dotted line) with its standard deviation  (shaded grey band). Open diamonds represent the observed host galaxies of \\citet{Sav09}, while filled circles correspond to their mean values in redshift intervals of 0.5. For comparison, we also include the corresponding mean trends for the low metallicty sample (dashed line) and the complete galaxy population (solid line).} \\label{gsfr} \\end{figure}   Regarding the mean specific SFR (SSFR), defined as the ratio between the SFR and the stellar mass of a given galaxy (Fig.~\\ref{gssfr}), host galaxies are predicted to show SSFRs similar to those in the low metallicity sample, and higher than those in the complete galaxy population, for $z < 2$. Between $z \\sim 2$ and $z \\sim 6$, there are no significant differences in mean  SSFRs among the three samples, which have larger SSFRs that low redshift galaxies. This can be understood because, at higher redshift, galaxies have larger gas reservoirs which can feed stronger star formation activity and are, on average, less chemically enriched. At low redshift, host galaxies seem to be particularly efficient at transforming gas into stars. Note, however, that the mean SSFR predicted by scenario II.3 at low redshift is half an order of magnitude lower than the mean observed SSFR.  \\begin{figure} \\includegraphics[width=0.45\\textwidth]{fig4.ps} \\caption{Mean  SSFR of the  host galaxies as a function of redshift predicted by scenario II.3 (dotted line) with its standard deviation  (shaded grey band). Open diamonds represent the observed host galaxies of \\citet{Sav09}, while filled circles correspond to their mean values in redshift intervals of 0.5. For comparison, we also include the corresponding mean trends for the low metallicty sample (dashed line) and the complete galaxy population (solid line).} \\label{gssfr} \\end{figure}   \\subsection{Luminosity and colour}  Observations show that host galaxies tend to be bluer than the general population of galaxies observed at a given redshift, and fainter than a typical $L^*$ galaxy. In fact, this trend can be nicely reproduced by our scenario II.3, as shown in Fig.~\\ref{colours}. The  luminosities and colours of the predicted host galaxies are in excellent agreement with the observations compiled by \\citet{Sav09}. From this figure, we can see that host galaxies are bluer than the complete galaxy population for $ z < 2$, but have similar mean colours for higher redshfits. Our scenario predicts that host galaxies are more luminous systems in the $B$-band, compared to the mean luminosity of the global galaxy population at all redshifts.  \\begin{figure} \\includegraphics[width=0.45\\textwidth]{fig5_1.ps} \\includegraphics[width=0.45\\textwidth]{fig5_2.ps} \\caption{Mean $B$-band luminosity (upper panel) and $B-K$ colour (lower panel) as a function of redshift predicted by scenario II.3 (dotted line) with its standard deviation (shaded grey band). Open diamonds represent the  observed host galaxies  by \\citet{Sav09}, while filled circles correspond to their mean values in redshift intervals of 0.5. For comparison, we also include the corresponding mean trends for the low metallicty sample (dashed line) and the complete galaxy population (solid line).} \\label{colours} \\end{figure}  \\subsection{Metallicity}  In Fig.~\\ref{metal}, we compare the mean cold gas metallicity of the host galaxies predicted by scenario II.3 to different metallicity estimations reported by \\citet{Sav09}. Since at $z > 0.2$ it is very difficult to distinguish HII regions with enough resolution, these authors generally measured the optical luminosity-weighted mean metallicity in a galaxy. For consistency, host galaxies at $z<0.2$ were treated by them as the rest of the sample (i.e. integrating fluxes over the whole galaxy). In this figure, we have included metallicities obtained by using different indicators \\citep[see][for more details]{Sav09}. In some cases, absorption lines in the optical afterglow can be used to obtain the metallicity of neutral cold gas along the line of sight of the GRB (the so-called GRB-DLAs). This is the case of 9 GRB-DLA systems studied by \\citet{Sav06}, all of them at $z > 1.6$. In this case, the metallicity could be associated more directly to the metallicity of the host galaxy, contrary to QSO-DLAs which are associated to HI clouds in the intergalactic medium. At $z<1$, measurements of metallicity by \\citet{Sav09} are derived from hot gas and, in this case, the lower branch solution is preferred for the hosts. As shown in Fig.~\\ref{metal}, the predicted host galaxies metallicities are systematically higher than the GRB-DLA metallicities. This might indicate that LGRBs occur in regions of even lower metallicity than the mean metallicity of the cold gas of the host galaxies. In this sense, \\citet{Nuz07}, adopting similar hypotheses to generate synthetic LGRB populations, but using full cosmological simulations where the chemical enrichment of baryons was consistently followed with redshift, claimed a metallicity threshold of $0.3 Z_{\\odot}$ for the progenitor stars in order to reproduce observations. More detailed information on the metallicity of individual stellar populations, which also takes into account the inhomogeneities of the interstellar medium, are needed to improve our understanding of this issue (see \\citealt{Pon09} and Artale et~al. in preparation). Also along these lines, a work by \\citet{Mod08} compares the metallicity of GRB hosts associated to Type Ic supernovae (SN Ic) to broad-line SN Ic without GRB detection at $z < 0.25$. The metallicity of SN-GRB host galaxies are estimated by computing the nebular oxygen abundance by three different methods. In this case, they directly probe the environments of the progenitors because either they were measured at that location or they correspond to metallicities of homogeneous dwarf galaxies. \\citet{Mod08} results are in agreement with the values obtained by \\citet{Sav09} and with our own results at low redshifts.  \\begin{figure} \\includegraphics[width=0.45\\textwidth]{fig6.ps} \\caption{Mean cold gas metallicity as a function of redshift predicted by scenario II.3 (dotted line) with its standard deviation (shaded grey band). Filled circles represent observed metallicities of host galaxies reported by \\citet{Sav09}, measured using the $T_e$ technique. We also include those measured by the R32 upper branch (triangles) or lower branch (diamonds), and GRB-DLAs (asterisks). For comparison, we include the corresponding mean trends for the low metallicty sample (dashed line) and the complete galaxy population (solid line).} \\label{metal} \\end{figure}   \\section[]{Environment of LGRB hosts} \\label{env}  It is still under discussion if host galaxies are always located in regions of a given characteristic local density or if, depending on $z$, they  inhabit different environments. Recently, \\citet{Cam09} estimated the cross-correlation function of their simulated host galaxies finding that these systems tend to map underdense regions. Observations provide somehow contradictory results. \\citet{Bor04} found host galaxies to inhabit field regions. Lately, \\citet{Wai07} suggested that host galaxies might be biased towards interacting or merging galaxies. It is yet too soon for observations to provide a robust answer to these questions but theoretical predictions could positively contribute to this area. The advantage of our model is that taking into account the observability allows us to make a more robust prediction.  In this section, we investigate the environment of the simulated host galaxies by means of three different estimators. The distance to the closest neighbour ($d_1$) of a host galaxy provides an estimation of the possibility of a host galaxy being interacting with another galaxy. The distance to the fifth neighbour ($d_5$) of a host galaxy provides an estimation of the density of the region the galaxy inhabits. Finally, the central halo virial mass ($M_{\\rm vir}$) is an estimator of the global potential well to which the host galaxy is bound.  In the particular case of the closest neighbour, it is well known that pair interactions can enhance the star formation activity \\citep[e.g.][]{Lam03,Pat05} and if LGRB progenitors are massive stars this would imply an increase in the probability of detecting a LGRB event. Since semi-analytical codes do not model tidally-induced star formation, we modified the SFR for galaxies with their closest neighbour located at distances smaller than $0.1 h^{-1}\\,{\\rm Mpc}$ by a factor of two, motivated by observational and numerical results \\citep[e.g.][]{Lam03,Per06,Per09,DiM08}. However, this correction does not change any of the trends reported in this work. But the role of interactions should be further tested with models that properly include their effects.  We estimate the cumulative distribution of these three environmental estimators ($d_1,d_5, M_{\\rm vir} $) predicted by scenario II.3. To quantify behaviours as a function of redshift, we determine the values $d_{1,50},d_{5,50}$ and $M_{\\rm vir,80}$ at which the corresponding cumulative fractions reach $50$ per cent in the case of $d_1$ and $d_5$, and $80$ per cent in the case of $ M_{\\rm vir}$. In Fig.~\\ref{d1} (upper panel), we show $d_{1,50}$ as a function of redshift according to scenario II.3, together with the trend for the low-metallicity sample and the complete galaxy population. At high redshift, the three samples tend to have the nearest neighbour at approximately similar distances. We only note a slight trend for host galaxies to systematically have a closer first neighbour than galaxies in the other two samples for $z > 1$ . From $z \\sim 1$, low metallicity galaxies tend to have their closest neighbour further away than either a galaxy in the complete population or a host galaxy. The latter tend to inhabit lower density regions than galaxies in the complete population. Our models indicate that observable host galaxies would only have a slightly higher probability to be in a galaxy pair compared with the general galaxy population at high redshift (i.e. $50$ per cent  of galaxies have the first neighbour closer than $\\sim 50\\,{\\rm kpc}$).  The global environment is quantified by $d_{5,50}$. As shown in Fig.~\\ref{d1} (middle panel) at high redshift ($z > 2$), galaxies in the three samples reside in similar environments. For lower redshift there is a clear trend for galaxies in the low metallicity sample and host galaxies to inhabit lower density regions than galaxies in the complete galaxy population. This results agree with the numerical work of \\citet{Cam09}, and also with the observational findings of \\citet{Bor04}, taking into account that the mean redshift of LGRBs observed at that time was $\\langle z \\rangle < 2$.  In the case of the virial mass, as shown in Fig.~\\ref{d1} (lower panel), we find that  $80$ per cent  of observable host galaxies have haloes less massive than $10^{11.0-11.5}\\, {\\rm M_\\odot}$ at any redshift. As the structure forms and groups and clusters aggregate hierarchically, galaxies tend to inhabit larger dark matter haloes. However, host galaxies stay within a narrower range of mass haloes. The inversion in the relation observed in this figure can be understood considering that host galaxies observability depends strongly on the star formation activity which, in turn, depends strongly on environment \\citep{Pog09}. Hence, from $z \\sim 2$ active star forming galaxies (i.e. those which are important generators of LGRBs) with mean stellar masses of $\\sim 10^{10}\\, {\\rm M_{\\odot}}$  reside again in slightly smaller dark matter haloes (Fig.~\\ref{massred} and Fig.~\\ref{gsfr}). Note that this is not the case for the general low metallicity sample which although residing, on average, in smaller haloes than the global galaxy population, tend to systematically inhabit larger ones with decreasing redshift, as expected in a hierarchical scenario. The observability condition which is closely linked to the star formation activity produces this kind of halo downsizing scenario for host galaxies at low redshift, while the mean stellar mass remains approximately constant.    \\begin{figure} \\includegraphics[width=0.45\\textwidth]{fig7_1.ps} \\includegraphics[width=0.45\\textwidth]{fig7_2.ps} \\includegraphics[width=0.45\\textwidth]{fig7_3.ps} \\caption{Environmental study of host galaxies predicted by scenario II.3 (diamonds), the complete galaxy population (stars), and the low metallicity galaxy sample (crosses). {\\em Upper panel:} Median distance to the closest neighbour ($d_{1,50}$). {\\em Middle panel:} Median distance to the fifth neighbour ($d_{5,50}$). {\\em Left panel:} Virial mass of the central halo at the $80$  percentile ($M_{\\rm vir,80}$).} \\label{d1} \\end{figure} ",
        "conclusions": "\\label{con}  Our results indicate that the observed LGRB host galaxies properties can be reproduced assuming that LGRB progenitors are massive stars  and occur in galaxies with moderately low mean cold gas metallicities. The joint requirement for the synthetic LGRBs to reproduce the observations of BATSE and for the simulated  systems that host them  to reproduce the  stellar mass distribution of the observed host galaxies determines the minimum mass for the progenitor star ($m > 75 M_{\\odot}$)\\footnote{The mass threshold for LGRB progenitors was computed for the \\citet{Cha03} IMF used in the catalogue of \\citet{DeL07}, and would be  54 M$_{\\odot}$ if a \\citet{Sal55} IMF had been used, instead.} and the maximum metallicity cut-off ($0.6\\,{\\rm  Z_{\\odot}}$) for the cold gas metallicity of the simulated host galaxies. Our scenario II.3, which satisfies these conditions, succeeds at reproducing the dependence on redshift of stellar mass, luminosity, colour, SFR, SSFR and metallicity of the observed host galaxies compiled by \\citet{Sav09} over the redshift range $0 < z < 3$.   Our main findings are:  \\begin{enumerate}  \\item The average stellar mass of host galaxies is higher than the average stellar mass of the complete galaxy population, and remains within a narrow mass range around $10^{10}\\,{\\rm M_\\odot}$ , with a very weak trend to lower mass for higher redshifts,  due to the double requirement of low mean gas metallicity and high star formation rates. We note that the  dispersion is approximately 0.5 dex around the $\\log$ mean stellar mass.  \\item Compared to the characteristic stellar mass $M^* \\sim 10^{10.63}\\,{\\rm M_{\\odot}}$ estimated by \\citet{Bel03}, host galaxies are low stellar mass systems (88 per cent of our host galaxies have stellar masses between $\\sim 10^{8.5-10.3}\\, {\\rm M_\\odot}$). However, compared to the mean stellar mass of the complete galaxy population, host galaxies tend to be massive galaxies.  \\item The SFR of host galaxies is larger than that of galaxies in the complete population at all redshifts. At $z > 2$, the SSFR of observable host galaxies is comparable to the complete galaxy population, but at lower redshift it is systematically higher. As a consequence, host galaxies are, on average, bluer than the global galaxy population. The  dispersion found for  the properties of the observable host galaxies reflect the different histories of formation of galaxies at a given stellar mass.  \\item Our results support the claims for a metallicity threshold to reproduce the properties of the current observed host galaxies. The global metallicity threshold of $0.6\\, {\\rm Z_{\\odot}}$ we derived from our models is an upper limit for the metallicity of the LGRB progenitor stars in the case of no dust effects. The comparison of the mean metallicity of host galaxies predicted by our model with GRB-DLAs observations suggests that LGRBs might be produced in stars of even lower metallicity as pointed out by other authors \\citep[e.g.][]{Nuz07,Cam09}. At low redshift, our results are in agreement with SN-GRB local metallicity estimations by \\citet{Mod08}. However, observations show a large spread in metallicity and there are also observational uncertainties that can affect its determination. Hence, results should be taken with caution.  \\item For $1< z < 6$, host galaxies seem to be slightly more likely to be in pairs than galaxies belonging to the other samples. There is, however, a change of behaviour for $z < 1$, where the at least 50 per cent of the host galaxies seems to have their closest neighbour further away than galaxies in the complete population, but closer than galaxies in the low metallicity sample. We highlight that we are considering neighbouring galaxies of a mass above $10^9\\,M_{\\odot}$. This limit is set by  the numerical resolution of the {\\em Millennium Simulation}. Lower mass companions could also imprint morphological perturbations and trigger star formation activity \\citep[e.g.][]{Lam03} but we cannot pursue this analysis further on with this galaxy catalogue.  \\item Regarding global environment, our model suggests that,  at $z > 1$, observable  host galaxies would preferentially inhabit environments of similar density to those populated by the general population  and of slightly higher density than those inhabited by the low metallicity sample. Towards $z = 0$, observable host galaxies tend to be  progressively located in less dense environments, which becomes as subdense as the regions where low metallicity galaxies reside \\citep[see also][]{Cam09}.   \\item Our results suggest that  observable host galaxies tend to have dark matter haloes in the range $10^{11.0-11.5}\\,M_{\\odot}$, regardless of redshift, and show a slight signal for halo downsizing from $z \\sim 1$, distinguishing them from the other two samples which follow the expected halo mass growth in a hierarchical scenario. This result is consistent with observable host galaxies being systems with  mean stellar masses approximately constant, regardless of the age of the Universe.  \\end{enumerate}  Our results are mainly a consequence of the joint requirements to have high star formation activity to ensure observability and to reproduce the current distribution of stellar masses of observed host galaxies. However, within the current constrains, galaxies with high masses would be too metal-rich to produce  LGRBs and low mass systems would have low probability of being observed. If the observed  sample were modified by the incorporation of other galaxies, for example dusty hosts, then our model would need to be readjusted accordingly."
    },
    "1005/1005.1878_arXiv.txt": {
        "abstract": "An exact solution describing the evolution of the type Bang-to-Rip with the phantom divide line crossing is constructed in the Chameleon cosmology model, based on two independent functions of the scalar field. ",
        "introduction": "The discovery of cosmic acceleration\\cite{cosmic} has stimulated an intensive study of  models of dark energy\\cite{dark} responsible for the origin of this phenomenon. Dark energy is characterized by a negative pressure whose relation to the energy density $w=p/\\varepsilon$ is less than $-1/3$. Moreover, if this relation happens to be less than $-1$ such kind of dark energy is called ``phantom'' dark energy\\cite{phantom}. The breakdown of the condition $w > -1$  implies the so called super-acceleration of cosmological evolution, which in some models culminates approaching a new type of cosmological singularity called Big Rip\\cite{Rip}. While the cosmological constant ($w = -1$) is still a possible candidate for the role of dark energy, there are observations giving  indications in favor of the models where the equation of state parameter $w$ not only changes with time, but could become less than $-1$ no\\cite{observe}. There are also observations \\cite{obs-cross} indicating the phenomenon of the phantom divide line crossing \\cite{divide}, when the parameter $w$ changes between the phantom and non-phantom phases. A standard way to introduce the phantom energy is to consider a scalar  field with the negative sign of the kinetic energy term, usually called ``phantom'' scalar field. The simplest way of describing the phenomenon of the phantom divide line crossing is to introduce two scalar fields - phantom and non-phantom as in the so called quintom models \\cite{two-exp}.  Some other models are based on the use of the PT symmetry ideas \\cite{we-PT,we-PT1,we-PT2}. Note that the phantom divide line crossing can occur also in models based on one scalar field if some special potential are chosen \\cite{we} or if the non-minimal coupling between a scalar field and gravity is included \\cite{non-minimal}.  Naturally, consideration of two scalar fields open wider opportunities for description of any prescribed cosmological evolution, than working with a single field \\cite{we-two}.  Recently, another class of cosmological dark energy models was put forward - the so called chameleon cosmology models \\cite{chameleon}. In this models the scalar field interacts with some kind of matter, behaving as perfect fluid. The Lagrangian of the model contains a term, where the matter interacts with the effective metric (physical metric, multiplied by a conformal factor depending on the scalar field). The energy density of the perfect fluid is proportional to the trace of its energy-momentum tensor, multiplied by this conformal factor. On the other hand the effective mass of the scalar field becomes variable, depending on the energy density of matter, which changes in the process of the cosmological evolution.  Lately, an alternative  version of the Chameleon cosmology was proposed in paper \\cite{Fara-Sale}. Here the term, describing the interaction between the chameleon scalar field and matter has a very general multiplicative form and is equal to some arbitrary function of the scalar field multiplied by the Lagrangian of  matter, which on the Friedmann-Robertson-Walker background coincides with the pressure of the matter. Thus, one has two arbitrary functions of the scalar field to play with - the standard potential and the potential-like function describing the interaction of matter with the chameleon field. Choosing properly  these two functions of the scalar field one can describe different kinds of cosmological evolution, including those where the universe undergoes the phantom divide line crossing. The authors of paper \\cite{Fara-Sale} use the numerical integration of the corresponding equations of motion and compare the results with the data coming from time evolution of the cosmological redshift of distant sources.  We remark that having two functions of a scalar field one gets a sufficient freedom to construct exactly solvable solutions of the cosmological equations, giving some prescribed cosmological evolution. Just like choosing exponential potentials one can find exact solutions desribing power-law expansion of the universe (see e.g. \\cite{power-law}), here one can try to construct two potential-like functions providing the existence of exact cosmological solutions, undergoing the phantom divide line crossing. In fact, this is the task of the present paper.  We consider the flat Friedmann model with the metric \\begin{equation} ds^2 = dt^2 - a^2(t)dl^2. \\label{Friedmann} \\end{equation} The Hubble parameter behaves as \\begin{equation} H(t) \\equiv \\frac{\\dot{a}(t)}{a(t)} = \\frac{\\alpha t_{R}}{t(t_R-t)}, \\label{Hubble} \\end{equation} where $\\alpha$ is some positive parameter of the order $1$ and ``dot'' means as usual the derivative with respect to the cosmic time $t$. Such a cosmological evolution was considered in papers \\cite{we-PT1,we-two} in the the context of two-field models and was called ``Bang-to-Rip'' evolution. Here the universe begins its evolution from the Big Bang singularity when the Hubble parameter behaves as $H(t) \\sim \\frac{\\alpha}{t}$ and ends in the Big Rip singularity at the moment $t = t_R$. At the moment $t = \\frac{t_R}{2}$ the phantom divide line is crossed by the universe. Indeed, the phantom divide line crossing means that the relation between the pressure and the energy density becomes equal to $-1$. That means that in the energy conservation law $\\dot{\\varepsilon} = -3H(\\varepsilon+p)$ the right-hand side vanishes, and so does the time derivative of the total energy density in its left-hand side. The energy density is proportinal in the flat Friedmann model to the Hubble parameter squared. Hence, the time derivative of the Hubble parameter $H$ should vanish too. The time derivative of the Hubble parameter from Eq. (\\ref{Hubble}) is \\begin{equation} \\dot{H} = \\frac{2\\alpha t_R(2t-t_R)}{t^2(t_R-t)^2} \\label{Hubble1} \\end{equation} and vanishes at $t = t_R/2$.   We shall consider the  minimally coupled to gravity scalar field $\\phi$ with the potential $V(\\phi)$ whose interaction with the perfect fluid is described by the term in the  Lagrangian, which on the Friedmann background looks like \\cite{Fara-Sale}: \\begin{equation} L_{scalar+matter} = \\gamma\\rho f(\\phi), \\label{Cham} \\end{equation} where the coefficient $\\gamma$ relates the energy density $\\rho$ and the pressure $p$ of matter: \\begin{equation} p = \\gamma \\rho, \\label{eq-of-state} \\end{equation} and $f(\\phi)$ is some function of the scalar field $\\phi$. In the next section we shall find the form of the functions $V(\\phi)$ and $f(\\phi)$, and the exact solution for $a(t)$ and $\\phi(t)$ matching the evolution (\\ref{Hubble}). In the third section we discuss the bounds on the parameters of the model. The last section contains some concluding remarks. ",
        "conclusions": "We have constructed an exact  solution for a particular Chameleon cosmological model of the type considered in paper \\cite{Fara-Sale}. The universe in this solution begins its evolution from Big Bang singularity, undergoes a phantom divide line crossing and ends in the Big Rip singularity. The two potential-like functions of the chameleon scalar field have a rather  simple analytic form. Note that this form is simpler than the potential functions in two-scalar model, providing the same cosmological evolution \\cite{we-PT1}.  Thus, we hope that our version of the Chameleon model can be useful for the analysis of the phantom divide line crossing phenomenon."
    },
    "1005/1005.1563_arXiv.txt": {
        "abstract": "It is shown that monotonic depletion of the magnetic field pressure in  a quaking neutron star undergoing Lorentz-force-driven torsional seismic vibrations about axis of its dipole magnetic moment is accompanied by the loss of vibration energy of the star that causes its vibration period to lengthen at a rate proportional to the rate of  magnetic field decay.  Highlighted is the magnetic-field-decay induced conversion of the energy of differentially rotational Alfv\\'en vibrations into the energy of oscillating magneto-dipole radiation. A set of representative examples of magnetic field decay illustrating the vibration energy powered emission with elongating periods produced by quaking neutron star are considered and discussed in the context of theory of magnetars. ",
        "introduction": "Most, if not all,  reported up to now computations of frequency spectra of poloidal and toroidal Alfv\\'en vibration modes in pulsars and magnetars rest on tacitly adopted assumption about constant-in-time magnetic field in which a perfectly conducting neutron star matter undergoes Lorentz-force-driven oscillations (e.g. Carroll et al. 1986, Bastrukov et al. 1997, 1999, Lee 2007, 2008, Shaisultanov, Eichler 2009, Sotani et al. 2009; see, also, references therein). A special place in the study of the above Alfv\\'en modes of pure shear magneto-mechanical seismic vibrations ($a$-modes) occupies a homogeneous model of a solid star with the uniform density $\\rho$ and frozen-in poloidal static magnetic field of both homogeneous and inhomogeneous internal and dipolar external configuration. This field can be conveniently represented in the form \\begin{eqnarray} \\label{e1.1} {\\bf B}({\\bf r})=B\\,{\\bf b}({\\bf r}),\\quad {\\bf b}({\\bf r})=[b_r({\\bf r})\\neq 0, b_\\theta({\\bf r})\\neq 0, b_\\phi({\\bf r})=0],\\quad B={\\rm constant}, \\end{eqnarray} where $B$ is the field intensity [in Gauss] and ${\\bf b}({\\bf r})$ stands for the dimensionless vector-function of spatial distribution of the field.  In particular, the above form of ${\\bf B}({\\bf r})$ has been utilized in recent works (Bastrukov et al. 2009a, 2009b, 2009c) in which the discrete spectra of frequencies of node-free torsional Alfv\\'en oscillations has been computed in analytic form on the basis of equations of linear magneto-solid mechanics (Bastrukov et al. 2009a) \\begin{eqnarray} \\label{e1.1a} && \\rho{\\ddot {\\bf u}}=\\frac{1}{c} [\\delta {\\bf j}\\times {\\bf B}],\\quad \\nabla\\cdot {\\bf u}=0,\\\\ \\label{e1.1b} && \\delta {\\bf j}=\\frac{c}{4\\pi}[\\nabla\\times \\delta {\\bf B}], \\quad \\delta {\\bf B}=\\nabla\\times [{\\bf u} \\times {\\bf B}],\\quad \\nabla\\cdot {\\bf B}=0. \\end{eqnarray} Here ${\\bf u}={\\bf u}({\\bf r},t)$ is the field of material displacement, the fundamental dynamical variable of solid-mechanical theory of elastically deformable (non-flowing) material continua.  These equations describe  Lorentz-force-driven non-compressional vibrations of a perfectly conducting elastic matter of a solid star with Amp\\'ere form of  fluctuating current density $\\delta {\\bf j}$. Equation for $\\delta {\\bf B}$ describing coupling between fluctuating field of material displacements ${\\bf u}$ and background magnetic field ${\\bf B}$ pervading stellar material is the mathematical form of Alfv\\'en theorem about frozen-in lines of magnetic field in perfectly conducting matter\\footnote{Equation (\\ref{e1.1a}) can be represented in the following equivalent tensor form (Bastrukov et al. 2009a), ${\\ddot u}_{i}=\\nabla_k \\tau_{ik}$, where stress tensor $\\tau_{ik}=(1/4\\pi)[B_i\\delta B_k+B_k\\delta B_i-B_j\\delta B_j\\delta_{ik}]$ is the Maxwellian tensor of  magnetic field stresses. This last equation for $u_i$ is identical in appearance to  canonical equation of solid-mechanics ${\\ddot  u}_{i}=\\nabla_k \\sigma_{ik}$ where $\\sigma_{ik}=2\\mu\\,u_{ik}+[\\kappa-(2/3)\\mu]\\,u_{jj}\\delta_{ik}$ is the Hookean stress tensor and $u_{ik}=(1/2) [\\nabla_{i}\\,u_k+\\nabla_k\\,u_i]$ is the tensor of shear deformations in an isotropic elastic continuous matter with shear modulus $\\mu$ and bulk modulus $\\kappa$. The adopted for Alfv\\'en vibrations of compact solid stars terminology [which is not new of course, e.g., Moon (1984)], namely, magneto-solid mechanics  and solid-magnetics, is used as a solid-mechanical counterpart  of well-known terms like magneto-fluid mechanics, magnetohydrodynamics (MHD) and hydromagnetics, $\\rho \\delta {\\dot {\\bf v}}=(1/c) [\\delta {\\bf j}\\times {\\bf B}],\\quad \\delta {\\bf j}=(c/4\\pi)[\\nabla\\times \\delta {\\bf B}],\\quad \\delta {\\dot {\\bf B}}=\\nabla\\times [\\delta {\\bf v} \\times {\\bf B}]$, describing flowing magento-active plasma in terms of the velocity of fluctuating flow $\\delta {\\bf v}={\\dot {\\bf u}}$  and fluctuating magnetic field ${\\delta {\\bf B}}$. The MHD approach is normally utilized in astrophysics of main-sequence (MS) liquid stars. The advantage and virtue of the Rayleigh energy method of computing frequencies of node-free Alfv\\'en vibrations (Bastrukov et al. 2009a, 2009b, 2009c, Molodtsova et al. 2010) is that it can too be applied to the above MHD equations underlying  the study of $a$-modes in liquid MS-stars such, for instance,  as rapidly oscillating Ap (roAp) stars, chemically peculiar magnetic stars exhibiting high-frequency oscillations that have been and still remain the subject of extensive investigation (e.g., Ledoux and Walraven 1958, Rincon and Rieutord 2003).}. As for the general asteroseismology of compact objects is concerned, the above equations  seems to be appropriate not only for neutron stars but also white dwarfs (Molodtsova et al. 2010) and quark stars (Heyvaerts et al. 2009). The superdense material of these latter yet hypothetical compact stars is too expected to be in solid state (Xu 2003, 2009), that is, in the aggregate state possessing mechanical property of shear elasticity.  Inserting (\\ref{e1.1a}) in (\\ref{e1.1b}) the former equation of solid-magnetics takes the form \\begin{eqnarray} \\label{e1.2} && \\rho\\, {\\ddot {\\bf u}}({\\bf r},t)=\\frac{1}{4\\pi} [\\nabla\\times[\\nabla\\times [{\\bf u}({\\bf r},t)\\times {\\bf B}({\\bf r})]]]\\times {\\bf B} ({\\bf r}). \\end{eqnarray} The obtained equation serves as a basis of our further analysis. The studied in above works vibration regime is of some interest in that the rate of differentially rotational material displacements \\begin{eqnarray} \\label{e1.3} && {\\dot {\\bf u}}({\\bf r},t)= [\\mbox{\\boldmath $\\omega$}({\\bf r},t)\\times {\\bf r}],\\quad \\mbox{\\boldmath $\\omega$}({\\bf r},t)=[\\nabla\\chi({\\bf r})]\\,{\\dot\\alpha}(t), \\\\ \\label{e1.4} && \\nabla^2\\chi({\\bf r})=0,\\quad \\chi({\\bf r})=f_\\ell\\,P_\\ell(\\cos\\theta), \\quad f_\\ell(r)=A_\\ell\\, r^{\\ell} \\end{eqnarray} has one and the same form as in torsion elastic mode of node-free vibrations under the action of Hooke's force of mechanical shear stresses\\footnote{It may be worth noting that in standing-wave regime of elastic-force-driven differentially rotational vibrations, the toroidal field of material displacements is described one and the same equation (\\ref{e1.3}), but function $\\chi=A_\\ell j_\\ell(kr) P_\\ell(\\cos\\theta)$, is the solution of Helmholtz equation $\\nabla^2\\chi+k^2\\chi=0$ where $k^2=\\omega^2/c^2_t$ and $c_t^2=\\mu/\\rho$ is the speed of shear wave of material displacements in solid stellar matter of density $\\rho$ and whose shear modulus is $\\mu$. The discrete frequency spectrum $\\omega(\\ell)=2\\pi\\nu(\\ell)$ of elastic torsional vibration of multipole degree $\\ell$ is computed from from boundary condition $n_k\\sigma_{ik}=0$ resulting in dispersion equation, $dj_\\ell(z)/dz=j_\\ell(z)/z$, where $z=kR$. The characteristic feature of the long-wave regime of node-free vibrations, $z<<1$, with the toroidal field of material displacements (\\ref{e1.3})-(\\ref{e1.4}) is the absence of dipole overtone in the spectrum of discrete frequencies (Bastrukov 2007a, 2007b): $\\nu_e(\\ell)=\\nu_e[(2\\ell+3) (\\ell-1)]^{1/2}$, with basic frequency of shear elastic vibrations, $\\nu_e=c_t/R$, where $R$ stands for the star radius. In the non-studied before  regime of short-wave length limit of torsional  elastic vibrations (the limit $z>>1$ taken in the exact dispersion equation for dipole overtone $dj_1(z)/dz=j_1(z)/z$) leads to the spectrum of standing-wave dipole vibrations of the following analytic form $\\nu_e(\\ell=1)=\\nu_e(p+1)$, where $p=0,1,2,3...$ is the node number of short-wavelength vibration regime.} (Bastrukov et al. 1999, 2007, 2008). Hereafter $P_\\ell(\\cos\\theta)$ stands for Legendre polynomial of degree $\\ell$ specifying the overtone of $a$-mode. The time-dependent amplitude $\\alpha(t)$ describes temporal evolution of above vibrations; the governing equation for $\\alpha(t)$ is obtained form equation (\\ref{e1.2}).  The prime purpose of above works was to get some insight into difference between spectra of discrete frequencies of toroidal $a$-modes in neutron star models having one and the same mass  $M$ and radius $R$, but different shapes of constant-in-time poloidal magnetic fields. By use of the energy method, it was found that each specific form of spatial configuration of static magnetic field  about axis of which the neutron star matter undergoes nodeless torsional oscillations is uniquely reflected in the discrete frequency spectra by form of dependence of frequency upon overtone $\\ell$ of nodeless vibration.   \\begin{figure*}[ht] \\begin{center} \\includegraphics[width=10cm]{f1.eps} \\end{center} \\caption{ The basic frequency and period of global Alfv\\'en oscillations, equations (\\ref{e1.6}), as functions of magnetic field intensity in the neutron star models with indicated mass and radius.} \\label{fig:fr-period} \\end{figure*}   The assumption about constant in time undisturbed magnetic field means that the internal magnetic field pressure, $P_B$, the velocity $v_A$ of Alfv\\'en wave in the star bulk (e.g. Lee 2007) \\begin{eqnarray} \\label{e1.5} P_B=\\frac{B^2}{8\\pi}, \\quad v_A=\\sqrt{\\frac{2P_B}{\\rho}} \\end{eqnarray} and, hence, the frequency $\\nu_A=\\omega_A/2\\pi$ (where $\\omega_A=v_A/R$) and the period $P_A=\\nu_A^{-1}$ of global Alfv\\'en oscillations \\begin{eqnarray} \\label{e1.6} \\nu_A=\\frac{B}{2\\pi}\\sqrt{\\frac{R}{3M}},\\quad P_A=\\frac{2\\pi}{B}\\sqrt{\\frac{3M}{R}} \\end{eqnarray} remain constant in the process of vibrations whose amplitude $\\alpha(t)$ subjects to standard equation of undamped harmonic oscillator (e.g., Bastrukov et al 2009a). The allow for viscosity of stellar material leads to exponential damping of amplitude, but the frequency $\\nu_A$ and, hence, the period $P_A$ preserve one and the same values as in the case of non-viscous vibrations (Bastrukov et al. 2009c). In Fig.1 these latter quantities are plotted as functions of intensity $B$ of undisturbed poloidal magnetic field in the neutron star models with indicated mass  $M$ and radius $R$.  The practical usefulness of chosen logarithmic scale in this figure is that it shows absolute vales of $\\nu_A$ and $P_A$ for global vibrations of typical in mass and radius neutron stars.  In this work, continuing a line of investigations reported in (Bastrukov et al. 2009a, 2009b, 2009c, Molodtsova et al. 2010),  we relax the assumption about constant-in-time magnetic field and examine the impact of its decay on the  vibration energy and period. In section 2, a mathematical background of quaking neutron star model is briefly outlined with emphasis on the loss of vibration energy caused by depletion of internal magnetic field pressure  and resulting vibration-energy powered magneto-dipole radiation. The decreasing of magnetic field pressure in the star is presumed to be caused by coupling between vibrating star and outgoing material which is expelled by quake, but mechanisms of star-envelope interaction resulting in the decay of magnetic field, during the time of vibrational relaxation, are not considered. It is shown that physically meaningful inferences regarding radiative activity of quaking neutron star can be made even when detailed mechanisms of depletion of magnetic field pressure in the process of vibrations triggered by quakes are not exactly known. In section 3, a set of representative examples illustrating the newly disclosed effects is considered. The summary of developed theory with emphasis on its relevance to electromagnetic activity of quaking magnetars is presented in section 4. ",
        "conclusions": "The main purpose of this work was to examine inferences of asteroseismic model of neutron star undergoing quake-induced torsional Alfv\\'en vibrations about axis of its dipole magnetic moment which are accompanied by monotonic depletion of internal magnetic field pressure. It is shown this latter process results in the loss of vibration energy and changes the character of temporal evolution of vibrations from harmonic in time vibrations (with constant frequency) in static, time-independent, magnetic field to non-isochronal vibrations with monotonically decreasing frequency in decaying magnetic field. The presented analysis suggests that the loss of vibration energy caused by depletion of magnetic field pressure should be accompanied by coherent (non-thermal) oscillating magneto-dipole radiation (with frequency of luminosity oscillations equal to the frequency of Alfv\\'en seismic vibrations of the star).  In this context it is appropriate to remind a seminal work of Hoyle, Narlikar and Wheeler (1964) in which it has been pointed out for the first time that vibrating neutron star should operate like Hertzian magnetic dipole deriving radiative power of magneto-dipole emission from the energy of magneto-mechanical vibrations (see, also, Pacini 2008). What is newly disclosed here is that conversion of vibration energy into the energy of magneto-dipole electromagnetic radiation can be realized when, and only when, torsional Alfv\\'en vibrations are accompanied by decay of magnetic field.  All above suggests that considered for the first time theory of vibration-energy powered emission of neutron star is relevant to electromagnetic activity of magnetars - neutron stars endowed with magnetic field of extremely high intensity the radiative activity of which is ultimately related to the magnetic field decay  (Duncan \\& Thompson 1992, Thompson \\& Duncan 1995). This subclass of highly magnetized compact objects is commonly associated with soft gamma repeaters and anomalous X-ray pulsars (e.g., Harding 1999, Kouveliotou 1999, Woods and Thompson 2006,  Mereghetti 2008) --  young  isolated and seismically active neutron stars (Chen et al. 1996, Franco et al. 2000). The magnetar quakes are exhibited by short-duration thermonuclear gamma-ray flash followed by rapidly oscillating X-ray flare of several-hundred-seconds duration. During this latter stage of quake-induced radiation  of magnetar, a long-periodic (5-12 sec) modulation of brightness was observed lasting for about 3-4 minutes (e.g., Harding 1999, Kouveliotou 1999). Taking into account such long-periodic and rapidly dying variability of magnetar pulsating radiation is the feature which is non-typical for young neutron stars powered by energy of rotation, it has been suggested (Bastrukov et al. 2002)  that detected long-periodic pulsating emission is powered by  the energy of torsional magneto-elastic vibrations of magnetar triggered by quake.  In view of key role of ultra-strong magnetic field it is quite likely that quasi-periodic oscillations (QPOs) of outburst flux from SGR 1806-20 and SGR 1900+14 discovered in observations reported in (Israel et al. 2005, Watts \\&  Strohmayer 2006, Terasawa et al. 2006) are produced by torsional seismic vibrations predominately sustained by Lorentz force (Bastrukov et al 2009b, 2009c) which are accompanied, as was argued above, by monotonic decay of background magnetic field. If so, the predicted elongation of QPOs period  of oscillating outburst emission from quaking magnetars should be traced in existing and future observations.  The authors are indebted to colleagues from Physics Department and Kavli Institute of Astronomy and Astrophysics, Peking University, Beijing, for valuable discussions and suggestions. Special thanks go to Alexey Kudryashov (Saratov University, Russia) who paid our attention to the existing of exact solution of obtained equation for vibration amplitude."
    },
    "1005/1005.5176_arXiv.txt": {
        "abstract": "We present deep, high-quality $K$-band images of complete subsamples of powerful radio and sub-millimetre galaxies at redshifts $z \\simeq 2$. The data were obtained in the very best available seeing via queue-based observations at the United Kingdom InfraRed Telescope and Gemini North, with integration times scaled to ensure that comparable rest-frame surface brightness levels are reached for all of the galaxies. We fit two-dimensional axi-symmetric galaxy models to these images to determine basic galaxy morphological parameters at rest-frame optical wavelengths $\\lambda_{rest} > 4000$\\AA, varying luminosity, axial ratio, half-light radius $r_{1/2}$, and S\\'{e}rsic index $n$. We find that, while a minority of the images show some evidence of galaxy interactions, $>95$\\% of the rest-frame optical light in all the galaxies is well-described by these simple axi-symmetric models. We also find evidence for a clear difference in morphology between these two classes of galaxy; fits to the individual images and to the image stacks reveal that the radio galaxies are moderately large ($\\langle r_{1/2} \\rangle = 8.4 \\pm 1.1$\\,kpc; median $r_{1/2}=7.8$ ), de Vaucouleurs spheroids ($\\langle n \\rangle = 4.07 \\pm 0.27$; median $n = 3.87$), while the sub-millimetre galaxies appear to be moderately compact ($\\langle r_{1/2} \\rangle = 3.4 \\pm 0.3$\\,kpc; median $r_{1/2}=3.1$\\,kpc) exponential discs ($\\langle n \\rangle = 1.44 \\pm 0.16$; median $n = 1.08$). Model fits to the stacked images yield very similar results. We show that the $z \\simeq 2$ radio galaxies display a well-defined Kormendy ($\\mu_e - r_e$) relation but that, while larger than other recently-studied massive galaxy populations at comparable redshifts, they are still a factor $\\simeq 1.5$ times smaller than their local counterparts. The scalelengths of the starlight in the sub-millimetre galaxies are comparable to those reported for the molecular gas, suggesting that the two may be co-located. Their sizes are also similar to those of comparably massive quiescent galaxies at $z > 1.5$, allowing the possibility of an evolutionary connection following cessation/quenching of the observed star-formation activity. In terms of stellar mass surface density, the majority of the radio galaxies lie within the locus defined by local ellipticals of comparable stellar mass. In contrast, while best modelled as discs at the epoch of bright dust-enshrouded star formation, most of the sub-millimetre galaxies have higher stellar mass densities than local galaxies, and appear destined to evolve into present-day massive ellipticals. ",
        "introduction": "\\label{intro} The statistical and individual properties of galaxies selected at different wavelengths, and at different redshifts must be explained in a self-consistent manner by any successful model of galaxy formation and evolution.  Driven by multi-wavelength surveys of ever increasing scale, recent years have seen substantial advances in our understanding of the basic demographics of different galaxy populations such as Lyman-break galaxies (e.g. Bouwens et al. 2007), radio galaxies (e.g. Rigby et al. 2010) and sub-millimetre galaxies (e.g. Coppin et al. 2006; Weiss et al. 2009; Austermann et al. 2010).  The resulting improving information in, for example, number counts and redshift distributions, has been fed back to modify/refine models of galaxy formation. However, important uncertainties and degeneracies in the basic physics generally remain. For example, sub-millimetre galaxies in particular were not a natural prediction of early semi-analytic models (Baugh et al. 2005), and there remains controversy over whether this can be fixed by the adoption of a top-heavy stellar intial mass function (IMF), or whether this tension is symptomatic of a basic inability of current semi-analytic models to produce enough star-formation activity in the young universe (Bower et al. 2006; Swinbank et al. 2008; Cirasuolo \\& Dunlop 2008).  To remove such uncertainties requires higher-level information. Specifically, different galaxy formation models make rather different predictions for galaxy masses, sizes and morphologies. For example, the models of Baugh et al. (2005) predict that sub-millimetre galaxies result from mergers of sub-$L*$ galaxies (with their apparent star-formation rates boosted by a top-heavy IMF). In contrast, the recent cosmological hydrodynamic simulations of Dav\\'{e} et al. (2010) predict that sub-millimetre galaxies are massive galaxies fed by smooth infall, are not typically involved in major mergers, and are simply the most extreme examples of normal star-forming galaxies (see also Finlator et al. 2006; Fardal et al. 2007; Dekel et al. 2009).  The estimation of galaxy stellar masses at high reshift has therefore become a major industry in recent years, boosted by the provision of rest-frame near-infrared photometry provided by the {\\it Spitzer Space Telescope}. The sensitivity of mass-to-light ratio on star-formation history is minimized at such wavelengths, but uncertainties due to choice of IMF or the relative importance of asymptotic giant branch stars have motivated attempts to obtain independent confirmation of galaxy masses at $z \\simeq 2$ via velocity dispersion measurements derived from emission-line spectroscopy at both millimetre and infrared wavelengths (e.g. Taconni et al. 2006, 2008; van Dokkum, Kriek \\& Franx 2009). While acknowledging the uncertainties, the available photometric and dynamical evidence suggests that, by the epoch of peak sub-millimetre emission at $z \\simeq 2-3$, sub-millimetre galaxies were already massive galaxies with estimated stellar  masses $M_{star} \\simeq 10^{11} M_{\\odot}$ (Greve et al. 2004; Tacconi et al. 2006; Borys et al. 2005; Clements et al 2008; Dye et al. 2008; Schael et al. 2010). This suggests that sub-millimetre galaxies represent an important phase in the formation of massive galaxies, and the challenge now is to establish their role in cosmic star-formation history, and the nature of their present-day descendents. Here we attempt to move this subject forward by measuring the basic sizes and morphologies of sub-millimetre galaxies at $z \\simeq 2$, performing such measurements at the longest possible wavelength (2.2\\,$\\mu m$) in order to trace the mass dominant stellar population, and to minimize the potentially confusing effects of dust obscuration at rest-frame ultraviolet wavelengths.  The measurement of galaxy sizes and morphologies at redshifts $z > 1$ has also been the subject of many studies in recent years, and has yielded some surprising and still controversial results. Specifically, a number of studies have now found that massive, passive galaxies at $1.5 < z < 3$ are much more compact than present-day galaxies of comparable mass (Trujillo et al. 2006, 2007; Longhetti et al. 2007; Cimatti et al. 2008; van Dokkum et al. 2008). The resulting inferred high stellar mass densities have been taken to imply that such objects cannot continue to evolve by purely passive evolution into any known type of present-day galaxy. These results are consistent with those of Toft et al. (2007) and Zirm et al. (2007) who, with the aid of {\\it Spitzer} MIPS photometry, were in addition able to show that, while passive high-redshift galaxies are surprisingly compact, star-forming galaxies at high redshift appear morphologically similar to normal star-forming disk galaxies. Recently Buitrago et al. (2008), from a {\\it Hubble Space Telescope} (HST) NICMOS study of 82 massive galaxies in GOODS North concluded that, for a given stellar mass, disc-like galaxies at $z \\simeq 2.3$ were a factor 2.6 $\\pm$ 0.3 smaller than present-day systems, while spheroid-like galaxies were a factor 4.3 $\\pm$ 0.7 smaller. They also noted that this evolution is at least in qualitative agreement with some hierarchical semi-analytic predictions (e.g. Khochfar \\& Silk 2006). However, there is now growing evidence that not all high-mass galaxies at $z > 1$ follow this trend, with some at least appearing to be fully formed (Saracco, Longhetti \\& Gargiulo 2010).  At present it is unclear where sub-millimetre galaxies fit into this picture. As star-forming galaxies they might be expected to be extended discs, like Lyman-break galaxies and present-day star-forming galaxies. On the other hand, if they are the progenitors of today's massive ellipticals, their apparently high-density starbursts (Tacconi et al. 2006) might be related to the high stellar densities reported in red, passive galaxies at high redshift ($\\sim~2 \\times 10^{10} {\\rm M_{\\odot}\\, kpc^{-3}}$, comparable to present-day globular clusters).  To date, genuinely deep, $K$-band imaging has only been obtained for a small number of sub-millimetre galaxies (e.g. Smail et al. 2003), and generally with inadequate resolution for the reliable extraction of basic morphological parameters. Very recently, Swinbank et al. (2010) reported the results of HST NICMOS+ACS imaging of a sub-sample of sub-millimetre galaxies. However, they did not report reliable measurement of, for example S\\'{e}sic indices for individual galaxies, and the wavelength limit of HST prevented them from undertaking observations in the $K$-band, which ensures imaging longward of the 4000\\AA/Balmer break out to $z \\simeq 3$.  A high-quality, systematic, deep, $K$-band imaging study of a well defined sample of sub-millimetre galaxies is thus desirable, and has the potential to clarify several fundamental issues. In this paper we attempt to take a step towards better quantifying the sizes, morphologies and stellar mass densities of sub-millimetre galaxies via a deep, and high-quality $K$-band imaging study of a well-defined sample of luminous SMGs at $z \\simeq 2$. To aid the interpretation of our results we have taken the somewhat unusual step of obtaining images of equivalent depth and resolution for a complete comparison sample of the most luminous radio galaxies known at comparable redshifts ($z \\simeq 2$).  The radio galaxy sample is a useful control sample because, given the now-well-established proportionality of black hole and spheroid mass (e.g. McLure \\& Dunlop 2002), we can be very confident that these galaxies will end up as massive ellipticals at the present-day. However, given the above-mentioned evidence that many massive galaxies at $z > 1$ are more compact than comparably massive present-day objects, a study of radio galaxy sizes at $z \\simeq 2$ is also clearly of interest in its own right. Rather surprisingly, while deep near-infrared imaging of 3CR radio galaxies has been undertaken at $z \\simeq 1$ (e.g. Dunlop \\& Peacock 1993; Best et al. 1998) very deep, high-quality $K$-band imaging of the most powerful 3CR radio galaxies at $z \\simeq 2$ has not in fact been undertaken before (although NICMOS $H$-band images of a sample of even higher redshift radio galaxies were obtained by Penterrici et al. 2001). Consequently, while the morphologies of radio galaxies out to $z \\simeq 1$ are rather well defined (McLure \\& Dunlop 2000; McLure et al. 2004), their rest-frame optical morphologies and sizes at $z \\simeq 2$ have yet to be clarified.  To undertake this study we have used NIRI on Gemini North to image the sub-millimetre galaxies, and UFTI on UKIRT to image the radio galaxies. We have taken great care to ensure that all images are as well matched as possible in $K$-band surface-brightness limit, designing integration times to compensate for the difference in telescope aperture between Gemini North and UKIRT, and for cosmological surface-brightness dimming within our samples. Also crucial to this study has been our ability (due to the queue-based scheduling system in operation at both telescopes) to demand the very best seeing conditions for all of our observations ($FWHM < 0.6$ arcsec). It is only with this powerful combination of consistent high-quality seeing and long exposure times (typically 90 mins on Gemini) that it has proved possible to determine the sizes (i.e. half-light radii) and basic morphologies (i.e. S\\'{e}rsic index) of these massive $z \\simeq 2$ galaxies via ground-based imaging.  The paper is structured as follows. In Section 2 we summarize the radio galaxy and sub-millimetre galaxy samples, and the existing supporting information. In Section 3 we present the UKIRT and Gemini $K$-band imaging observations, and in Section 4 we explain how these data were reduced. Then in Section 5 we describe how we modelled these new images to extract the basic morphological properties of the galaxies. The results of the individual model fits are presented, summarized and discussed in Section 6, with the images, model fits and residual images provided in Appendix A. We also explore the robustness of our results, through simulations, non parametric tests of galaxy size, and the analysis of radio  galaxy and sub-millimetre galaxy image stacks. Finally, in Section 7 we attempt to place our results in the context of studies of other galaxy populations at both high and low redshift. Our main conculsions are summarized in Section 8.  \\begin{table*} \\begin{center}  \\caption[Details of the radio and sub-millimetre galaxy $K$-band observations]{Details of the radio and sub-millimetre galaxy $K$-band observations. Column 1 lists the IAU name for the radio galaxies, and the 8-mJy survey name for the sub-millimetre sources. Column 2 gives the original Cambridge and Parkes catalogue names for the radio sources, and the alternative SHADES catalogue names for the Lockman sub-millimetre galaxies (Coppin et al. 2006). Column 3 provides information on the nature of the galaxy identification for the sub-millimetre sources. Column 4 names the telescope with which the data were acquired, and Column 5 lists the total exposure time in seconds for each target. Column 6 provides the measured 6-arcsec diameter aperture $K$-band magnitude (in the Vega magnitude system) for each galaxy, along with the associated error. Finally, Columns 7, 8, and 9 list the adopted redshift for each source, the type of redshift, and the reference for the source from which this redshift was obtained: A from Best el al. (1999), B from Hewitt \\& Burbidge (1991), C from Spinrad et al. (1985), a from Aretxaga et al. (2005), b from Aretxaga, priv. comm., c from Chapman et al. (2003), d from Chapman et al. (2005), e from Ivison et al. (2005), f from Smail et al. (2003).}  \\begin{tabular}{llllrclcc} \\hline\\hline Source  &  Radio Catalogue Name  &   &  Telescope & Exp. time (sec)  &  $K$ mag. (Vega) & $z$  &  $z$-type & z-ref\\\\ \\hline 0016$-$12   &  3C 008\t     & & UKIRT &  9480    & \\phantom{$>$}18.13$\\pm$0.07 & 1.589      & $z_{spec}$            &  A\\\\ 0128$-$26   &  PKS 0128$-$26 & & UKIRT &  18840   & \\phantom{$>$}17.77$\\pm$0.09 & 2.348      & $z_{spec}$            &  A\\\\ 0231$+$31   &  3C 068.2      & & UKIRT &  9720    & \\phantom{$>$}17.97$\\pm$0.08 & 1.575      & $z_{spec}$            &  B\\\\ 0310$-$15   &  PKS 0310$-$15 & & UKIRT &  8100    & \\phantom{$>$}18.35$\\pm$0.11 & 1.769      & $z_{spec}$            &  A\\\\ 0851$-$14   &  PKS 0851$-$14 & & UKIRT &  10140   & \\phantom{$>$}17.87$\\pm$0.08 & 1.665      & $z_{spec}$            &  A\\\\ 1008$+$46   &  3C 239\t     & & UKIRT &  11820   & \\phantom{$>$}17.89$\\pm$0.07 & 1.790      & $z_{spec}$            &  C\\\\ 1019$+$22   &  3C 241        & & UKIRT &  9600    & \\phantom{$>$}17.75$\\pm$0.08 & 1.617      & $z_{spec}$            &  C\\\\ 1120$+$05   &  3C 257        & & UKIRT &  22560   & \\phantom{$>$}17.58$\\pm$0.07 & 2.474      & $z_{spec}$            &  B\\\\ 1140$-$11   &  PKS 1140$-$11 & & UKIRT &  13920   & \\phantom{$>$}18.79$\\pm$0.09 & 1.935      & $z_{spec}$            &  A\\\\ 1422$-$29   &  PKS 1422$-$29 & & UKIRT &  9600    & \\phantom{$>$}17.69$\\pm$0.07 & 1.632      & $z_{spec}$            &  A\\\\ 1533$+$55   &  3C 322\t     & & UKIRT &  9180    & \\phantom{$>$}17.86$\\pm$0.08 & 1.681      & $z_{spec}$            &  B\\\\ 1602$-$17   &  PKS 1602$-$17 & & UKIRT &  15120   & \\phantom{$>$}17.63$\\pm$0.09 & 2.043      & $z_{spec}$            &  A\\\\ 2356$+$43   &  3C 470        & & UKIRT &  10260   & \\phantom{$>$}18.11$\\pm$0.12 & 1.653      & $z_{spec}$            &  B\\\\ \\hline \\hline Source      &  SHADES Name   & ID Type     & Telescope & Exp. time (sec)&  $K$ mag. (Vega)    & $z$          & $z$-type & z-ref\\\\ \\hline N2 850.01   &                & Radio ID    & Gemini N &  4750    & \\phantom{$>$}19.90$\\pm$0.15 & 2.850      & $z_{spec}$ &  d\\\\ N2 850.02   &                & Radio ID    & Gemini N &  4500    & \\phantom{$>$}20.77$\\pm$0.27 & 2.454      & $z_{spec}$ &  d  \\\\ N2 850.04   &                & Radio ID    & Gemini N &  4625    & \\phantom{$>$}18.85$\\pm$0.10 & 2.387      & $z_{spec}$ &  f  \\\\ N2 850.06   &                & ERO/VRO     & Gemini N &  5500    & \\phantom{$>$}20.18$\\pm$0.16 & 3.0        & $z_{phot}$ &  b  \\\\ N2 850.07   &                & Radio ID    & Gemini N &  5000    & \\phantom{$>$}20.00$\\pm$0.16 & 3.1        & $z_{phot}$ &  b  \\\\ N2 850.08   &                & Radio ID    & Gemini N &  5000    & \\phantom{$>$}18.30$\\pm$0.09 & 1.190      & $z_{spec}$ &  d  \\\\ N2 850.09   &                & Radio ID    & Gemini N &  4750    & \\phantom{$>$}20.99$\\pm$0.37 & 3.3        & $z_{phot}$ &  b  \\\\ N2 850.12   &                & ERO/VRO     & Gemini N &  5500    & \\phantom{$>$}20.31$\\pm$0.17 & 2.5        & $z_{phot}$ &  a  \\\\ N2 850.13   &                & Radio ID    & Gemini N &  5375    & \\phantom{$>$}21.61$\\pm$0.37 & 2.283      & $z_{spec}$ &  d  \\\\ LE 850.02   & LOCK850.3      & Radio ID\t   & Gemini N &  5500    & \\phantom{$>$}19.66$\\pm$0.11 & 3.041      & $z_{spec}$ &  d  \\\\ LE 850.03   & LOCK850.17     & Radio ID\t   & Gemini N &  5500    & \\phantom{$>$}19.66$\\pm$0.12 & 2.239      & $z_{spec}$ &  d/e\\\\ LE 850.04   & LOCK850.6      & ERO/VRO     & Gemini N &   5500    & \\phantom{$>$}20.74$\\pm$0.14 & 3.6       & $z_{phot}$ &  b  \\\\ LE 850.06   & LOCK850.14     & Radio ID\t   & Gemini N &  5625    & \\phantom{$>$}19.16$\\pm$0.17 & 2.610      & $z_{spec}$ &  c/e\\\\ LE 850.07   & LOCK850.16     & Radio ID\t   & Gemini N &  5500    & \\phantom{$>$}18.87$\\pm$0.10 & 1.9        & $z_{phot}$ &  b  \\\\ LE 850.08   & LOCK850.41     & Radio ID\t   & Gemini N &  5500    & \\phantom{$>$}20.38$\\pm$0.16 & 3.4        & $z_{phot}$ &  b  \\\\ \\hline \\end{tabular} \\label{tabsample}  \\end{center} \\end{table*} ",
        "conclusions": "We have obtained deep, high-quality $K$-band images of complete subsamples of powerful radio and sub-millimetre galaxies at redshifts $z \\simeq 2$. The data were obtained in the very best available seeing (FWHM $\\simeq$ 0.5\\,arcsec) through the queue-based observing systems at the United Kingdom InfraRed Telescope and Gemini North, with integration times scaled to ensure that comparable rest-frame surface brightness levels are reached for all of the galaxies. We have analysed the resulting images by fitting two-dimensional axi-symmetric galaxy models to determine basic galaxy morphological parameters at rest-frame optical wavelengths $\\lambda_{rest} > 4000$\\AA, varying luminosity, half-light radius $r_{1/2}$, and S\\'{e}rsic index $n$. We have also undertaken an extensive set of simulations which show that, for galaxies with scalelengths $r_{1/2} > 2$\\, kpc, the modelling of our deep, high-quality ground-based images should yield half-light radii and S\\'{e}rsic indices accurate to $\\simeq 10$\\%.  We find that, while a minority of the images show some evidence of galaxy interactions, $>95$\\% of the rest-frame optical light in all the galaxies is well-described by these simple axi-symmetric models. We also find evidence for a clear difference in morphology between the radio galaxies and the sub-millimetre galaxies. Specifically, the fits to the individual images reveal that the radio galaxies are moderately large ($\\langle r_{1/2} \\rangle = 8.4 \\pm 1.1$\\,kpc; median $r_{1/2}=7.8$ ), de Vaucouleurs spheroids ($\\langle n \\rangle = 4.07 \\pm 0.27$; median $n = 3.87$), while the sub-millimetre galaxies appear to be moderately compact ($\\langle r_{1/2} \\rangle = 3.4 \\pm 0.3$\\,kpc; median $r_{1/2}=3.1$\\,kpc) exponential discs ($\\langle n \\rangle = 1.44 \\pm 0.16$; median $n = 1.08$). Model fits to the stacked images yield very similar results.  We find that the $z \\simeq 2$ radio galaxies display a well-defined Kormendy ($\\mu_e - r_e$) relation, which is offset in surface brightness from the local relation by $\\simeq 1.7$ mag. This luminosity offset is consistent with that expected due to purely passive evolution between $z \\simeq 2$ and $z \\simeq 0$ for a stellar population formed at $z > 3$. However, while larger than other recently-studied massive galaxy populations at comparable redshifts, the $z \\simeq 2$ radio galaxies are still on average a factor $\\simeq 2$ times smaller than their local counterparts, suggesting that they have still to undergo further dynamical evolution.  The scalelengths we have derived for the starlight in the sub-millimetre galaxies are consistent with those reported for the molecular gas (e.g. Tacconi et al. 2008), suggesting that the two may be co-located. Such scalelengths are also similar to those of comparably-massive quiescent galaxies at $z > 1.5$, allowing the possibility of an evolutionary connection following cessation/quenching of the observed star-formation activity.  Finally, in terms of stellar mass surface density, we find that the majority of the radio galaxies lie within the locus defined by local ellipticals of comparable stellar mass. In contrast, while best modelled as discs at the epoch of bright dust-enshrouded star formation, most of the sub-millimetre galaxies have higher stellar mass densities than local galaxies, and appear destined to evolve into present-day massive ellipticals.  With the advent of {\\it Herschel} and SCUBA-2, coupled with the wide-field near-infrared imaging provided by VISTA and WFC3/IR on {\\it HST}, it should soon be possible to extend this type of study to larger samples of sub-millimetre galaxies which include even more extreme starbursts at $z \\simeq 2$. It will then be possible to test whether the dense disc morphologies of sub-millimetre galaxies uncovered in the present study are a universal feature of dust-enshrouded starbursts at high redshift, or whether they are specific to the mass/star-formation rate regime sampled in the early sub-millimetre surveys."
    },
    "1005/1005.5389_arXiv.txt": {
        "abstract": "\\noindent Noticeable deviations from the prediction of the fiducial LCDM cosmology are found in the angular power spectrum of the CMB. Besides large-angle anomalies, the WMAP $1^{st}$ year data revealed a dip in the power spectrum at $l \\sim 200$, which seemed to disappear in the $3^{rd}$ year and subsequent angular power spectra. Using the WMAP single $1^{st}$, $3^{rd}$, and $5^{th}$ year data as well as the total $5$ year coadded data, we study the intensity and spatial distribution of this feature in order to unveil its origin and its implications for the cosmological parameters.  We show that in all WMAP data releases there is a substantial suppression of the first Doppler peak in a region  near the north ecliptic pole. ",
        "introduction": "The dynamics of both individual galaxies and clusters of galaxies, the Hubble diagram for distant Type Ia supernovae and the pattern of temperature anisotropies in the Cosmic Microwave Background (CMB) broadly support the standard flat $\\Lambda$-dominated cosmological model with a nearly scale-invariant spectrum of adiabatic Gaussian primordial fluctuations, such as might be generated at the end of an inflationary epoch, as an accurate description of our universe.  In particular, this simple model provides an acceptable fit to the Wilkinson Microwave Anisotropy Probe (WMAP) observations of the CMB temperature anisotropies. The pattern of acoustic peaks and troughs predicted by the model to be imprinted on the angular power spectrum by the primordial fluctuations  in the inflationary field and its subsequent evolution has been detected and characterized to good precision~\\cite{ Hinshaw:2003ex,Spergel:2003cb,Hinshaw:2006ia}. A global fit to the data has enabled cosmologists to determine the handful of parameters of the inflationary $\\Lambda$CDM model, which does a remarkably good job at describing the spectrum.  The WMAP data, however, also exhibits unexpected anomalies that have sparked considerable attention \\cite{Bennett:2010jb,Copi:2010na}. These include excesses or deficiencies of power in at least three bins -- $\\ell \\sim 22$, $44$ and $200$ --  violations of statistical isotropy at $\\ell<6$~\\cite{de OliveiraCosta:2003pu,Schwarz:2004gk,Land:2005ad}, and a severe  lack of large angle  correlations requiring the low-$\\ell$ $C_\\ell$ to not be independent~\\cite{Spergel:2003cb,Copi:2006tu}. Hemispheric asymmetries in the temperature power spectrum extending over a wide range of angular scales as well as hemispheric dependent non-gaussian signatures were reported in~\\cite{Eriksen:2003db,Hansen:2008ym,Rath:2007ti,Rossmanith:2009cy}, and a non-Gaussian cold spot in the southern galactic hemisphere was found using wavelet transforms~\\cite{Vielva:2003et,Cruz:2004ce}. Evidence for primordial non-Gaussianity of the local type in the temperature spectrum in the third year data has been claimed in~\\cite{Yadav:2007yy} (However, it does not seem to be present when later data are included in the analysis~\\cite{Komatsu:2008hk,Smith:2009jr}).  If these anomalies are of primordial origin, rather than caused by unaccounted foreground or secondary effects, there could be profound implications for our understanding of the early Universe. For example, while the simplest single field inflationary scenarios predict negligible deviations from Gaussianity, larger deviations can be accommodated in extended particle physics models~\\cite{Bartolo:2004if}. The lack of power in some low-$\\ell$ multipoles, though not the multipole alignments and not the lack of large-angle correlations, could be explained by the presence of anticorrelated isocurvature perturbations such as those expected in the curvaton model~\\cite{Gordon:2003fi,Ferrer:2004nv}. Topological defects could be at the origin of the  cold spot~\\cite{Cruz:2007pe}.  Most of these anomalies were first noted in the first year WMAP data release and persisted into the subsequent data releases~\\cite{Copi:2006tu,Eriksen:2007pc}. The dip in the first peak found in the first year WMAP data release is believed to be a notable exception -- it is no longer present when the newly adopted noise weighting scheme  is used to extract the best fit power spectrum. This anomaly in the first year angular power spectrum was, thus, attributed to a noise fluctuation~\\cite{Hinshaw:2006ia}.  Surprisingly, the dip in the power spectrum around $\\ell \\sim 200$ disappears when data from the {\\em ecliptic} poles is not included in the analysis of the 1st year WMAP data regardless of the weighting scheme (see Fig. 7 in~\\cite{Hinshaw:2003ex}). Archeops,  the sole previous ground-based experiment that observed the region of the sky around the  north ecliptic pole, showed a dip around the first peak. All other similar experiments showed no such dip and had access only to southern ecliptic regions~\\cite{Wang:2002rta}. This suggests that the differences might not simply be due to a noise fluctuation, but might rather reflect a real anomaly.  In this paper we take a closer look at the degree-scale CMB temperature anisotropies. We estimate the primordial power spectrum  using data from different regions of the sky according to the two weighting schemes that have been used by the WMAP collaboration. Our findings show that there is a substantial suppression of the first Doppler peak in a region near the north ecliptic pole. This reduction in the temperature power spectrum around the first acoustic peak occurs regardless of the weighting scheme used, although the detailed shape (i.e. the dip like feature) is only present when the analysis follows the strategy used by WMAP in their first year release and might also depend on the particular binning used to present the data.  The rest of the paper is organized as follows. In section~\\ref{sec:method}, we review the method used for estimating the power spectrum, $C_{\\ell}$, from observations. Section~\\ref{sec:data} describes the datasets that we used, together with the processing steps taken to minimize the influence of non-cosmological foregrounds. In section~\\ref{sec:stat} we discuss the method used to determine the statistical significance of our analysis of the data. Next, in section~\\ref{sec:results}, before our conclusions, we present the results of our analysis and briefly comment on implications for the physics underlying the CMB. ",
        "conclusions": "\\label{sec:conclusion} In the first year WMAP TT power spectrum, the binned $C_{\\ell}$ value at the first peak was significantly lower than the best-fit $\\Lambda$CDM value \\cite{Hinshaw:2003ex}. This was widely referred to as the ``dip in the first peak.\" In the third year data release the so-called dip disappeared. This was attributed to the use of an improved weighting scheme \\cite{Hinshaw:2006ia}. However, there was also shown to be a decrease in power at the first peak from the part of the sky near the ecliptic poles compared to the ecliptic plane in the first year data \\cite{Hinshaw:2003ex}. Indeed, using the improved (third year) weighting scheme, we have found that the ecliptic polar power at the first peak is significantly lower than the full-sky power in the first, third, fifth and coadded 5 year data releases. (We have not yet analyzed the seventh year data release, although we do not expect much difference.)  By separately masking out the north ecliptic polar region, the south ecliptic polar region and the ecliptic plane we have demonstrated that this difference can be ascribed to a large reduction in peak power in the north ecliptic polar region. We have shown that there is a significant increase in power when we mask out the north ecliptic pole for all weighting schemes and data releases. There is no such change in power when the south ecliptic pole is masked. Additionally, when using the third year weighting scheme there is still a sizeable reduction in power\twhen only the data from the ecliptic poles is used compared to when the full sky is analyzed.  We have characterized the absence of power at the first doppler peak in the temperature power spectrum in the north ecliptic polar region using the peak statistic defined in Eq. \\ref{eq:peak_stat}. Using either the first or third year weighting scheme and across all data releases the significance is at least 95\\% based on $10^{4}$ best-fit $\\Lambda$CDM realizations and the analysis of one $V \\times W$ cross-band power spectrum. The most trusted result should arguably be the significance of the coadded map with the third year weight scheme, as it provides analysis with the least noise and data processing method closest to what is currently used. Even when analyzing the most optimal data set for characterizing this effect, our result is still $\\approx 2\\sigma$ significant.  Indeed, the first peak may seem like a specific choice of range in $\\ell$ to study -- we would like to extend this analysis in the future to other ``special\" regions in $\\ell$-space, namely to the second and third doppler peaks. However current experiments are limited in their measurements around these peaks, so an accurate characterization of this effect is not possible. Future CMB experiments should provide better data for further analysis.  Archeops's observation of a dip in the first peak from measurements including the north ecliptic polar region of the sky, albeit at low statistical significance, suggests that this is not a WMAP systematic. Planck should therefore confirm the existence of a dip in the first peak power near the north ecliptic pole.  If these variations are of cosmological origin, they cannot be explained within the standard $\\Lambda$CDM model, since it predicts an isotropic and Gaussian spectrum. However, the association with the north ecliptic pole argues against a cosmological origin and in favor of Solar System physics. Possible signatures include a deviation in the CMB spectrum in this region of the sky and an associated anomaly in polarization maps."
    },
    "1005/1005.0423_arXiv.txt": {
        "abstract": "We examine the constraints imposed by helioseismic data on the solar heavy element abundances. In prior work we argued that the measured depth of the surface convection zone $R_{CZ}$ and the surface helium abundance $Y_{surf} $ were good metallicity indicators which placed separable constraints on light metals (CNONe) and the heavier species with good relative meteoritic abundances. The resulting interiors-based abundance scale was higher than some published studies based on 3D model atmospheres at a highly significant level.  In this paper we explore the usage of the solar sound speed in the radiative interior as an additional diagnostic, and find that it is sensitive to changes in the Ne/O ratio even for models constructed to have the same $R_{CZ}$ and $Y_{surf}$. Three distinct helioseismic tests (opacity in the radiative core, ionization in the convection zone, and the core mean molecular weight) yield consistent results. Our preferred O, Ne and Fe abundances are $8.86\\pm0.04$, $8.15\\pm0.17$ and $7.50\\pm0.05$ respectively. They are consistent with the midrange of recently published 3D atmospheric abundances measurements. The values for O, Ne and Fe which combine interiors and atmospheric inferences are $8.83\\pm0.04 $, $8.08\\pm0.09 $ and $7.49\\pm0.04 $ respectively. ",
        "introduction": "The Sun is a fundamental benchmark in astronomical abundance studies. Meteorites can be used to reliably infer relative abundances for numerous heavy elements and we have detailed information about limb darkening, precise fundamental parameters, and a wide variety of diagnostic features.  A decade ago the problem of solar abundances seemed settled. A series of papers \\citep[][GS98 hereafter]{ag89, gn93, gs98} had established a consistent set of surface abundances, and interiors models \\citep[e.g.][]{Bahcall2001} were in good agreement with helioseismic data ranging from sound speed as a function of depth to surface helium and convection zone depth measurements.  However, the atmospheres models used in this generation of study were one-dimensional and had important simplifying assumptions about the physics of the atmospheres.  With the advent of a new generation of numerical simulations it became possible to relax these severe assumptions.  For many heavier elements such as iron the revisions in abundance were small and within the errors.  However, the lighter elements (CNO) were a different story.  These species have problematic photospheric abundance diagnostics under solar conditions and meteorites cannot be used to infer their primordial relative abundGances. The 3D model atmospheres, along with other revisions such as enhanced non-LTE effects and the identification of blending features in commonly used abundance indicators,  led to much lower CNO abundances \\citep[][AGS05 hereafter]{ags05}.  The revised abundances placed solar interiors models in stark disagreement with helioseismic data \\citep[see, for example][]{AB05,BP04,BBPS05, Guzik05, Montalban04, Turck04}. This solar abundance problem spurred a re-examination of the opacities \\citep{Seaton04, Badnell05}, which improved opacity data as compared to the original source from the Opacity Project \\citep{SYMP1994, CJZ1995} but did not resolve the problem.  The next generation of helioseismic studies took a different approach, inverting the problem and using distinct characteristics to infer a higher seismic abundance more in line with the GS98 scale \\citep[][DP06 hereafter]{AB07, Chaplin07,DP06}. DP06 identified the depth of the surface convection zone and the surface helium abundance as sensitive abundance diagnostics.  This data could be used to predict absolute light and heavy metal abundances when combined with relative meteoritic abundances for heavy metals and relative photospheric/coronal abundances for CNONe. \\citet{AB07} used the depression of the adiabatic gradient induced by metal ionization, mostly from oxygen, to infer the absolute metallicty of the solar convection zone and obtained a comparably high abundance. \\citet{Chaplin07} derived a similar result from modelling the metallicity dependance of the mean molecular weight in the deep solar core.  There was a subsequent burst of revisions in the atmospheric abundances. The initial generation of models were criticised for problems in fitting the observed solar limb-darkening \\citep{ayres06, Koesterke08}.  Substantial revisions of the original model atmospheres led to a signficant improvement in the fit to solar data and removed the mild tension between the abundances of heavier elements (Fe, Si) and seismic inferences \\citep[][hereafter AGSS09]{AGSS09}.  However, the estimates for CNO abundances remained quite low.  Interestingly, a different group of investigators with an independant 3D atmospheres code had comparable agreement with solar data but found higher photospheric abundances \\citep{caffau08, caffau10}. The solar abundance problem therefore takes on a different cast; the distinction is not between new 3D and old 1D atmospheres, but rather on differing choices of abundance indicators and different ways of analyzing them.  A similar point was made by us in a prior paper \\citep{PD09}, where a reanalysis of the AGS05 work found an atmospheric abundance scale intermediate between seismology and the prior 3D values, comparable to the one obtained by \\citet{caffau10}.  However, neon has always been an important, and underconstrained, ingredient in the puzzle. Ne is a significant opacity source in stellar interiors, so the helioseismic problems induced by low oxygen could be compensated for if the solar Ne was sufficiently high \\citep{AB05,BBS05}. There are also chemical evolution arguments in favor of such a model \\citep{Turck04}. Ne is not preserved in meteorites in significant amounts and lacks photospheric diagnostics because of its atomic structure. As a result, indirect arguments from coronal abundance ratios must be used, and significant variations from the standard mixture have been argued for \\citep{drake05,drake07}. However, the bulk of evidence in the Sun appears to favor a lower ratio \\citep{Schmelz05, Young05}; see AGSS09 for a detailed discussion.  In DP06 we used only the scalar constraints ($R_{CZ} $ and $Y_{surf}$) as diagnostic features to infer abundances.  However, Ne retains electrons to higher temperatures than O does, and as a result different Ne+O mixtures with the same opacity at the base of the surface convection zone will have different opacity in deeper layers.  This leads to a distinct signature in the sound speed profile.  In this letter we update the mixture constraints based on the same method, but with improved stellar interiors model physics. By adding the sound speed profile as an additional constraint we can set interesting bounds on the solar $Ne/O$ ratio, and these are reinforced by considering the seismic metallicity bounds based on ionization and mean molecular weight. In section 2 we recall the method used in DP06, to determine the appropriate composition in order to reproduce the two observable scalars $R_{CZ}$ and $Y_{surf}$ and the extension of this method including the additional constraint $C_{sound}$ to determine the ratio $Ne/O$. In Section 3 we discuss our results. ",
        "conclusions": "Helioseismology has proven to be a powerful means of inferring the solar composition.  The depth of the surface convection zone and the initial helium abundance are very sensitive to the opacity of light and heavy metal respectively, and the errors in these model properties can be reliably quantified.  Reproducing these seismic scalar features require a relatively high solar metallicity.  If we combine these scalar constraints with reasonable priors about the relative abundances of CNONe (from the photosphere) and heavy elements (from meteorites) we can therefore infer a seismic solar mixture based primarily on stellar interiors physics. Recent revisions to the input physics of the solar model - improvements in the nuclear reaction cross-sections, low temperature opacities, and equation of state - yield a reference A(O) = 8.86 and A(Fe) = 7.50 which is the same as DP06.  However, the linkage between the thermal structure and the composition is indirect, which can lead to ambiguities in the interpretation of the strong constraints imposed by the measured solar thermal structure.  Fortunately, there is additional information encoded in the solar sound speed profile and there are independent helioseismic diagnostics which yield very similar results.  It is common to use differences in sound speed between models and the Sun as a measure of agreement. However, many of the striking visual features in such difference plots are in fact manifestations of errors in the independently determined scalar features, especially the convection zone depth.  We therefore argue that extracting information from the sound speed profile first requires calibration to a common set of observables, which can be achieved with modest adjustments of the solar composition.  Sound speed variations at the 0.1 - 0.2 \\% level arise from reasonable changes in the input physics. This sets a threshold below which deviations between the Sun and the models can be explained by known observational and theoretical sources.  It is possible to construct different relative mixtures of the light species CNONe which have similar opacities at the base of the convection zone, and thus similar scalar properties.  Changes in C and N are weakly constrained because their ionization structures are similar to that of oxygen and they are much less abundant (by the same token, even large variations in C and N cannot produce a low surface oxygen).  Neon, by contrast, is abundant enough that large variations could change the estimated oxygen, and uncertain enough that such variations cannot be dismissed out of hand.  We find that changes in the Ne/O level produce a strong differential sound speed signature at $R=0.56 R_{\\sun}$ which is explained by Ne retaining electrons (and thus higher opacity) to higher temperatures than O does.  Although the resulting bounds are modest (0.08 dex in oxygen), there would be a clear signal from a very high Ne / very low O mixture which are clearly not present in the data.  Our method is sensitive to opacity, and as a result the most likely systematic error source is in the theoretical quantum mechanical opacity calculations despite the very good agreement between independent theoretical calculation. Experiments are underway to check on this possibility and transition to an abundance scale primarily based on experiments.  However the same high abundances are inferred from two other independent helioseismic diagnostics - the level of metal ionization in the convection zone and the mean molecular weight of the solar core.  In a comprehensive review, \\citet{BA08} found that increases in neon did not reduce the need for a high oxygen in the convection zone (c.f. their Figure 21).  One would therefore need to invoke multiple errors in distinct interiors and atmospheric input physics, with the same sign and magnitude, to produce an acceptable high Ne - low O mixture.  Our best interiors estimate for composition is $A(O)=8.86\\pm0.04 $, $A(Ne)=8.15\\pm0.17 $ and $A(Fe)=7.50\\pm0.05$. \\citet{Lodders2009} combined recent atmospheres measurements for estimates of $A(O)=8.73\\pm0.07 $, $A(Ne)=8.05\\pm0.10 $ and $A(Fe)=7.45\\pm0.08$; these two scales are consistent within the errors. A weighted mean of the two leads to $A(O)=8.83\\pm0.04 $, $A(Ne)=8.08\\pm0.09 $ and $A(Fe)=7.49\\pm0.05$ which we contend is the most precise current estimate for abundances. In DP06 we derived relationships quantifying combinations of C, N, O, and Ne with the same $R_{CZ}$ and $Y_{surf}$. In the case of Ne, nonlinear effects can be induced when the change in the Ne/O ratio is too large and a better fitting relationship is $A(O)=  -7.52\\times (Ne/O)^4 +9.03\\times (Ne/O)^3 - 3.10\\times (Ne/O)^2 -0.32\\times (Ne/O) + 8.98$   Our abundance scale is consistent with some photospheric abundance measurements \\citep{caffau10,PD09} but not the lower AGSS09 values.  This is not a conflict between modern and primitive atmospheres treatments, or between one- and three-dimensional studies, but rather reflects differences between competing atmospheres models and judgment calls on the choice of indicators, continuum levels, and the proper treatment of blending features.  We are therefore hopeful that helioseismology may provide an absolute abundance standard which can be used to discriminate between competing models and which can be used to calibrate the appropriate composition diagnostics for the next generation of stellar models."
    },
    "1005/1005.5162_arXiv.txt": {
        "abstract": "\\noindent We present $H$- and $K_\\mathrm{s}$-band imaging data resolving the gap in the transitional disk around LkCa~15, revealing the surrounding nebulosity.  We detect sharp elliptical contours delimiting the nebulosity on the inside as well as the outside, consistent with the shape, size, ellipticity, and orientation of starlight reflected from the far-side disk wall, whereas the near-side wall is shielded from view by the disk's optically thick bulk.  We note that forward-scattering of starlight on the near-side disk surface could provide an alternate interpretation of the nebulosity. In either case, this discovery provides confirmation of the disk geometry that has been proposed to explain the spectral energy distributions (SED) of such systems, comprising an optically thick disk with an inner truncation radius of $\\sim$46\\,AU enclosing a largely evacuated gap. Our data show an offset of the nebulosity contours along the major axis, likely corresponding to a physical pericenter offset of the disk gap. This reinforces the leading theory that dynamical clearing by at least one orbiting body is the cause of the gap.  Based on evolutionary models, our high-contrast imagery imposes an upper limit of 21\\,$M_\\mathrm{Jup}$ on companions at separations outside of 0\\farcs1 and of 13\\,$M_\\mathrm{Jup}$ outside of 0\\farcs2.  Thus, we find that a planetary system around LkCa~15 is the most likely explanation for the disk architecture. ",
        "introduction": "The circumstellar disks of gas and dust around newly formed stars are believed to be the birthplaces of giant planets. In some protoplanetary disks, evidence of gaps or inner cavities has been revealed through analysis of the infrared spectral energy distribution \\citep[SED, e.g.][]{calvet2002} or interferometry at infrared \\citep[e.g.][]{ratzka2007} or millimeter wavelengths \\citep[e.g.][]{brown2009}.  These objects have been termed ``transitional disks'', since they are thought to represent a transitional state of partial disk dissipation between the protoplanetary disk stage and the debris disk stage, with a variety of mechanisms proposed to account for central cavities: Photoevaporation, disk instabilities, and dynamical clearing \\citep{alexander2006,chiang2007,bryden1999}.  One such system is \\object{LkCa 15} (K5:Ve, V=11.91\\,mag, J=9.42\\,mag, H=8.6\\,mag, K=8.16\\,mag), a 0.97\\,$M_\\odot$, 3--5\\,Myr old \\citep{simon2000}, weak-line T~Tauri star with a gas-rich millimeter-bright disk \\citep{pietu2007,henning2010}. Recently, it was identified as a pre-transitional disk \\citep{espaillat2007}, which consists of an inner disk component near the dust sublimation radius (0.12--0.15\\,AU), a wide gap, and an outer, disk beyond 46 AU \\citep{espaillat2008}.  Spitzer IRS spectra by \\citet{sargent2009} reveal a relatively low level of dust grain growth. With an expected angular scale of 0\\farcs33, the outer wall is accessible to 8-m class telescopes equipped with high-contrast imagers. ",
        "conclusions": ""
    },
    "1005/1005.4955_arXiv.txt": {
        "abstract": "{Most of what is known about clustered star formation to date comes from well studied star forming regions located relatively nearby, such as Rho-Ophiuchus, Serpens and Perseus. However, the recent discovery of infrared dark clouds may give new insights in our understanding of this dominant mode of star formation in the Galaxy. Though the exact role of infrared dark clouds in the formation process is still somewhat unclear, they seem to provide useful laboratories to study the very early stages of clustered star formation. Infrared dark clouds have been identified predominantly toward the bright inner parts of the galactic plane. The low background emission makes it more difficult to identify similar objects in mid-infrared absorption in the outer parts. This is unfortunate, because the outer Galaxy represents the only nearby region where we can study effects of different (external) conditions on the star formation process.} {The aim of this paper is to identify extended red regions in the outer galactic plane based on reddening of stars in the near-infrared. We argue that these regions appear reddened mainly due to extinction caused by molecular clouds and young stellar objects. The work presented here is used as a basis for identifying star forming regions and in particular the very early stages. An accompanying paper describes the cross-identification of the identified regions with existing data, uncovering more on the nature of the reddening.} {We use the Mann-Whitney U-test, in combination with a friends-of-friends algorithm, to identify extended reddened regions in the 2MASS all-sky $JHK$ survey. We process the data on a regular grid using two different resolutions, 60\\arcsec and 90\\arcsec. The two resolutions have been chosen because the stellar surface density varies between the crowded spiral arm regions and the sparsely populated galactic anti-center region.} {We identify 1320 extended red regions at the higher resolution and 1589 at the lower resolution run. The linear extent of the identified regions ranges from a few arc-minutes to about a degree.} {The majority of extended red regions are associated with major molecular cloud complexes, supporting our hypothesis that the reddening is mostly due to foreground clouds and embedded objects. The reliability of the identified regions is $>$99.9\\%. Because we choose to identify object with a high reliability we can not quantify the completeness of the list of regions.} ",
        "introduction": "\\label{intro} Dark clouds represent the earliest stages of star formation in the Galaxy. Nearby dark clouds have been studied in great detail and provide a wealth of information on the low-mass isolated star formation process \\citep[e.g.,][]{Shu1987}. However, most stars do not form in isolation but in groups or clusters \\citep{Lada2003}, and this clustered mode of star formation is not well understood. Some key questions of current star formation research are the origin of the stellar mass distribution, the relation between high-mass star formation and stellar clusters and the effect of environment on the star formation process. See \\citet{Zinnecker2007} for a recent review of this topic.  In the mid 1990s, infrared surveys with the Infrared Space Observatory (ISO) and the Midcourse Space Experiment (MSX) have revealed a class of dark clouds, referred to as Infrared Dark Clouds \\citep[IRDCs][]{Perault1996,Egan1998}. These clouds are observed in silhouette against the bright infrared background emission of the galactic plane. IRDCs are typically cold ($<$25\\,K) and dense ($\\sim$10$^5$\\,cm$^{-3}$) with masses ranging from \\citep[$\\sim$100--10$^5$\\,\\msun;][]{Egan1998,Carey1998,Carey2000} and they are located at larger distances ($>$1\\,kpc) compared to well-studied low-mass star forming regions. A recently published catalog of dark clouds extracted from the GLIMPSE Survey by \\citet{Peretto2009} shows that most IRDCs have column densities below N(H$_2$)$\\approx$5$\\times$10$^{22}$\\,cm$^{-2}$, with some extreme objects reaching N(H$_2$)$>$10$^{23}$\\,cm$^{-2}$. This new class of dark clouds are believed to represent the very early stages of clustered star formation, hence commonly referred to as cluster-forming clumps \\cite[e.g.,][]{Rathborne2006}.  While most of the star-forming activity in the Galaxy takes place in the inner spiral arms and the molecular ring, significant activity also occurs beyond the solar circle (R$_g$ $>$ 8.5\\,kpc). Prominent examples include well studied nearby ($<$500\\,pc) regions such as the Orion and Taurus-Perseus-Auriga complexes as well as more distant ($>$2\\,kpc) active star forming regions such as \\object{W3} and \\object{NGC 7538}. However, the outer Galaxy is generally somewhat neglected in studies of star formation compared with the inner Galaxy, which is unfortunate because it is a useful laboratory to study the effect of different external conditions on the process that initiates stellar birth. In particular, the metallicity, pressure, and radiation field in the outer Galaxy are considerably different from the inner Galaxy \\citep{Brand1995,Rudolph2006}, which makes it by far the most nearby region to study the effect of such variations on star formation.  The aim of this work is to locate star-forming regions in the outer galactic plane. To find concentrations of molecular gas in the outer galactic plane, where the mid-infrared background is insufficient to measure absorption, we use the extinction to background stars. The traditional method to do this is via star counts \\citep[e.g.,][]{Wolf1923}, which is limited by extinction to nearby clouds, especially when using optical data. This problem is avoided by using near-infrared data, but pick-up of emission from embedded objects makes this method unreliable.  An alternative to star counts is to use the near-infrared colour excess \\citep[NICE,][]{Lada1994} which is widely used to map the extinction toward dark clouds \\citep[e.g.,][]{Cambresy2002,Alves1998,Lada1999}. However, this method relies on knowledge of the distance to the object, so that a correction for foreground stars can be made. The presence of foreground stars can reduce the measured color excess, and hence may influence an identification based on this, significantly \\citep{Lombardi2005}. Furthermore, identifying objects in an automated, uniform way toward large regions in the sky is difficult using NICE because the extinction law is not uniform in every direction. Also, an average reference color of background stars is required, which is usually adopted from a nearby, well selected, extinction-free region.\\\\ \\ \\\\ The identification of dark clouds by the extinction they impose upon background stars requires a method that: \\begin{list}{}{\\topsep=5pt\\itemsep=0pt\\leftmargin=25pt\\rightmargin=15pt\\itemindent=0pt} \\item[1)]determines the extinction from the colors of stars, \\item[2)]can be applied to small numbers of stars (for high resolution), \\item[3)]can be treated uniformly in all directions, \\item[4)]puts a confidence level to measurements, \\item[5)]is independent of foreground stars. \\end{list} Because of the above mentioned limitations when using a near-infrared color excess method, we have chosen to search for candidate dark cloud regions in a statistical manner rather than based on the absolute value of the extinction.  In two papers, we describe the identification process of candidate regions in the outer galactic plane with the goal to build up a sample of objects that can be used for future studies of (clustered) star formation throughout the Galaxy.\\\\ \\ \\\\ In this first paper we describe the Mann-Whitney U-test. This distribution-free statistical test can be applied uniformly to large data-sets in order to identify parts of the data with deviating properties. The purpose of this paper is to identify extended red regions in the sky using data from the Two Micron All Sky Survey (2MASS). In an accompanying paper (Frieswijk et al. 2010, hereafter Paper II) we perform a cross-correlation between the red regions presented here and existing data (optical dark clouds, MSX, IRAS and CO where available). The purpose of the second paper is to provide an indication of the actual nature of the reddening toward the identified regions  In Section \\ref{data} we summarize the 2MASS data and discuss why we use the \\hk\\ colors for the identification. In Section 3 we explain the Mann-Whitney U-test and describe the procedure that we follow in order to determine the statistics. In Section 4 we describe the friends-of-friends approach that is used to extract extended red regions from the reliability images produced by the U-test. In Section 5 we present the results. A catalog of the extended red regions is available online\\footnote{\\\\ http://www.astro.rug.nl/$\\sim$ismgroup/OuterGalaxy/} and at the CDS. The reliability images of the entire outer galactic plane can be downloaded in {\\sc FITS} format from the same website. The final section summarizes our main conclusions.% ",
        "conclusions": "\\label{conclusions} We have used a statistical approach in combination with a friends-of-friends algorithm to measure deviations in the spatial \\hk\\ color distribution of stars in the outer galactic plane, with the goal to identify extended reddened regions. We processed the galactic plane at 60\\arcsec\\ and 90\\arcsec\\ resolution, resulting in the identification of 1320 high resolution and 1589 low resolution extended red regions. The reliability of individual red cells that make up each object is 99\\% and by using a friend-of-friends approach, the resulting catalog of objects is 99.9\\% reliable. Because we have put the effort in producing a highly reliable source list we cannot quantify the completeness of the catalog.  The majority of the objects are located toward well-known molecular cloud complexes and some correspond to specific, well-known objects such as \\object{NGC 7538} and \\object{S 159}. This correlation strengthens the argument that the red nature of the regions is caused by extinction and/or embedded young stellar objects rather than intrinsic star color variations. The main goal of our analysis is to find previously undetected dark clouds in the outer Galaxy. However, many other types of objects are included in the catalog and the results may be compared to other outer Galaxy studies in the future.  Based on the parameters derived from the statistical test it is impossible to determine the nature of the reddening. The next step is to cross-correlate the objects with existing optical, infrared and CO data. This process is described in an accompanying paper (Frieswijk et al. 2010)."
    },
    "1005/1005.3058_arXiv.txt": {
        "abstract": "{We present the dust properties of a small sample of Virgo cluster dwarf galaxies drawn from the science demonstration phase data set of the {\\em Herschel} Virgo Cluster Survey. These galaxies have low metallicities (7.8 $<$ 12 + log(O/H) $<$ 8.3) and star-formation rates $\\lesssim$ 10$^{-1}$ \\msun yr$^{-1}$. We measure the spectral energy distribution (SED) from 100 to 500 \\micron\\ and derive dust temperatures and dust masses. The SEDs are fitted by a cool component of temperature T $\\lesssim$ 20 K, implying dust masses around 10$^{5}$ \\msun and dust-to-gas ratios \\D  within the range 10$^{-3}$-10$^{-2}$. The completion of the full survey will yield a larger set of galaxies, which will provide more stringent constraints on the dust content of star-forming dwarf galaxies.} ",
        "introduction": "Late-type dwarf galaxies are metal-poor systems, and they represent unique environments to investigate the properties of dust in the low-metallicity regime. Mid/far-infrared (MIR/FIR) and submillimetre (submm) observations have provided information about the different dust components in dwarfs, showing that even metal-poor galaxies may host a significant amount of dust \\citep{1999ApJ...516..783T,2004ApJS..154..211H,2008ApJ...678..804E}. While the majority of the studies to date have focused mainly on bright and isolated dwarfs \\citep{2003A&A...407..159G,2005A&A...434..849H,2005A&A...434..867G,2009A&A...508..645G}, little is known about the interplay between the environment and the dust properties of low-mass and low-metallicity systems.  It is still unclear whether most of the dust mass in these systems is at very low temperature (T $\\lesssim$ 10 K), as the excess emission in the submm observed in some dwarfs might suggest \\citep{2002Ap&SS.281..247M,2005A&A...434..867G,2009A&A...508..645G}. This cold dust component may either be associated with the star-forming regions, residing in clumpy molecular complexes \\citep{2003A&A...407..159G}, or extend beyond the optical disc following the distribution of the neutral hydrogen \\citep[\\hi;][]{2002ApJ...567..221P}. However, the main objection to this interpretation is the high dust mass associated with the cold dust, and alternative explanations have been proposed \\citep[see][and references therein]{2010MNRAS.402.1409B}. For example, \\citet{2002A&A...382..860L}  used a mixed dust model \\citep{1990A&A...237..215D} of large and small grains, finding that a high abundance of small grains (compared to the Milky Way) with a shallow emissivity ($\\beta < 2$) may account for the excess submm emission in NGC\\,1569. \\citet{2009ApJ...706..941Z} applied a similar approach to  the starburst spiral galaxy NGC\\,3310  without introducing a cold dust component. Only additional measurements in the FIR/submm regime and the analysis of a larger sample of galaxies can help us properly assess the issue of dust at low temperatures.   The \\hers\\, Space Observatory \\citep{Pilbratt2010} allows us for the first time to  study the wavelength range between 200  and 500 \\micron, where the emission from cold dust is expected to be the most prominent. As part of the \\hers\\, Open Time Key Project, the \\hers\\, Virgo Cluster Survey\\footnote{http://www.hevics.org} \\citep[HeViCS;][]{Davies2010} will map an area of 64 square degrees of the Virgo cluster with PACS \\citep{Poglitsch2010} and SPIRE \\citep{Griffin2010} to investigate the dust content of the different morphological types within the cluster. The dwarf galaxy population of the Virgo cluster is dominated by dwarf ellipticals, but also contains a non-negligible fraction ($\\sim$ 10\\%) of late-type dwarfs with signs of current star formation activity \\citep{1987AJ.....94..251B}. In this Letter, we present the dust properties of three of these galaxies detected with PACS and SPIRE in the HeViCS Science Demonstration Phase (SDP) data set, a 4$^{\\circ} \\times 4^{\\circ}$ field covering the central region of the cluster.  \\begin{figure*} \\centering \\includegraphics[bb=14 14 300 300,width=3.cm,clip]{14653fg1a.eps} \\includegraphics[bb=54 360 270 576,width=3cm,clip]{14653fg1b.eps} \\includegraphics[bb=54 360 270 576,width=3cm,clip]{14653fg1c.eps} \\includegraphics[bb=54 360 270 576,width=3cm,clip]{14653fg1d.eps} \\includegraphics[bb=54 360 270 576,width=3cm,clip]{14653fg1e.eps} \\includegraphics[bb=54 360 270 576,width=3cm,clip]{14653fg1f.eps} \\\\ \\includegraphics[bb=14 14 300 300,width=3.cm,clip]{14653fg1g.eps} \\includegraphics[bb=54 360 270 576,width=3cm,clip]{14653fg1h.eps} \\includegraphics[bb=54 360 270 576,width=3cm,clip]{14653fg1i.eps} \\includegraphics[bb=54 360 270 576,width=3cm,clip]{14653fg1l.eps} \\includegraphics[bb=54 360 270 576,width=3cm,clip]{14653fg1m.eps} \\includegraphics[bb=54 360 270 576,width=3cm,clip]{14653fg1n.eps} \\\\ \\includegraphics[bb=14 14 300 300,width=3.cm,clip]{14653fg1o.eps} \\includegraphics[bb=54 360 270 576,width=3cm,clip]{14653fg1p.eps} \\includegraphics[bb=54 360 270 576,width=3cm,clip]{14653fg1q.eps} \\includegraphics[bb=54 360 270 576,width=3cm,clip]{14653fg1r.eps} \\includegraphics[bb=54 360 270 576,width=3cm,clip]{14653fg1s.eps} \\includegraphics[bb=54 360 270 576,width=3cm,clip]{14653fg1t.eps} \\caption{SDSS and \\hers \\, images of VCC\\,562, VCC\\,1179, and VCC\\,1356 at 100, 160, 250, 350, and 500 \\micron. The field size is 150\\arsec. The FWHM of the beam at each wavelength is displayed at the bottom-right corner. The different shape of the PACS beams depend on the single-scan (ellipse) or crossed-link scan (circle) coverage of the area.} \\label{fig:sdssspire}% \\end{figure*} ",
        "conclusions": "We have presented PACS and SPIRE observations of three star-forming dwarf galaxies in Virgo detected in the far-IR/submm regime with the science demonstration phase data set for the HeViCS survey. The data indicate the presence of cool dust with a temperature $\\la$20 K and dust masses around 10$^5$ \\msun. We have discussed the possibility that these galaxies host an additional cold (T $\\lesssim$ 10 K) component to explain the excess at 500 \\micron\\ in two of the dwarfs. However, the low S/N of the 500 \\micron\\ detections precludes us from drawing firm conclusions. The completion of the full area of the survey will enable us to place more stringent constraints on the dust content of star-forming dwarf galaxies in a dense cluster environment."
    },
    "1005/1005.4678_arXiv.txt": {
        "abstract": "\\PRE{\\vspace*{.3in}} The annihilation cross section of thermal relic dark matter determines both its relic density and indirect detection signals.  We determine how large indirect signals may be in scenarios with Sommerfeld-enhanced annihilation, subject to the constraint that the dark matter has the correct relic density. This work refines our previous analysis through detailed treatments of resonant Sommerfeld enhancement and the effect of Sommerfeld enhancement on freeze out. Sommerfeld enhancements raise many interesting issues in the freeze out calculation, and we find that the cutoff of resonant enhancement, the equilibration of force carriers, the temperature of kinetic decoupling, and the efficiency of self-interactions for preserving thermal velocity distributions all play a role.  These effects may have striking consequences; for example, for resonantly-enhanced Sommerfeld annihilation, dark matter freezes out but may then chemically {\\em recouple}, implying highly suppressed indirect signals, in contrast to naive expectations.  In the minimal scenario with standard astrophysical assumptions, and tuning all parameters to maximize the signal, we find that, for force-carrier mass $\\mphi = 250~\\mev$ and dark matter masses $m_X = 0.1$, 0.3, and 1 TeV, the maximal Sommerfeld enhancement factors are $\\Seff = 7$, 30, and 90, respectively.  Such boosts are too small to explain both the PAMELA and Fermi excesses.  Non-minimal models may require smaller boosts, but the bounds on $\\Seff$ could also be more stringent, and dedicated freeze out analyses are required.  For concreteness, we focus on $4\\mu$ final states, but we also discuss $4e$ and other modes, deviations from standard astrophysical assumptions and non-minimal particle physics models, and we outline the steps required to determine if such considerations may lead to a self-consistent explanation of the PAMELA or Fermi excesses. ",
        "introduction": "\\label{sec:intro}  Dark matter may be composed of thermal relics with mass near the weak scale $\\mweak \\sim 100~\\gev - 1~\\tev$.  Such candidates are produced in the hot early Universe and then freeze out when the Universe cools and expands.  Their annihilation cross section therefore determines both the relic density and the rate of annihilation today.  The requirement that the candidate be much or all of the dark matter therefore constrains its annihilation rate now, with important implications for indirect searches for dark matter.  Although annihilation in the early Universe and now is determined by the same dynamics, the kinematics are vastly different: at freeze out, thermal relics have relative velocity $\\vrel \\sim 0.3$, whereas today, $\\vrel \\sim 10^{-3}$.  As a result, the numerical values of the annihilation cross sections may differ significantly.  For example, neutralino annihilation is in many cases dominated by $P$-wave processes~\\cite{Goldberg:1983nd}, and so is suppressed at low velocities.  In this work, we consider Sommerfeld-enhanced cross sections, which have the opposite behavior: they are enhanced at low velocities and therefore boost present-day annihilation signals. We refine a previous study~\\cite{Feng:2009hw} and examine how large these boosts may be, subject to the constraint that the thermal relic has the correct thermal relic density to be dark matter.  Along with its general implications for future indirect searches, this study also has direct implications for the interpretation of current data. Following earlier excesses reported by the HEAT Collaboration~\\cite{Barwick:1997ig,Beatty:2004cy}, the PAMELA~\\cite{Adriani:2008zr}, ATIC~\\cite{:2008zzr}, and Fermi~\\cite{Abdo:2009zk} Collaborations have reported excesses of cosmic positrons over an estimate of expected background~\\cite{Strong:2009xj}.  Dark matter annihilation with Sommerfeld enhancement~\\cite{Sommerfeld:1931} of dark matter annihilation~\\cite{Baer:1998pg}, generalized to massive force carriers~\\cite{Hisano:2002fk,Hisano:2003ec,% Hisano:2004ds,Hisano:2005ec,Cirelli:2007xd}, has been proposed as an explanation~\\cite{Cirelli:2008pk,ArkaniHamed:2008qn,% Pospelov:2008jd,Fox:2008kb}.  In its simplest form, this scenario assumes that dark matter is composed of a single particle species $X$, which interacts with light force carriers $\\phi$ with fine structure constant $\\alpha_X$ and $\\mphi \\ll m_X \\sim \\mweak$.  This new interaction modifies dark matter annihilation and scattering properties.  The resulting annihilation cross section multiplied by relative velocity is $\\sv S$, where $\\sv$ is its tree-level value, and $S$ is Sommerfeld's original enhancement factor~\\cite{Sommerfeld:1931} \\begin{equation} S^0 = \\frac{\\pi \\, \\alpha_X / v} {1 - e^{- \\pi \\alpha_X / v}} \\ \\ \\stackrel{\\alpha_X \\gg v}{\\longrightarrow} \\ \\ \\frac{\\pi \\alpha_X}{v} \\ , \\label{S0} \\end{equation} generalized to massive $\\phi$, where $v \\equiv \\vrel/2$ is the dark matter particle's velocity in the center-of-mass frame.  In the proposed explanation, dark matter freezes out at early times when the Sommerfeld effect is negligible with $\\sv \\approx 3\\times 10^{-26}~\\cm^3/\\s$, leading to the correct relic density $\\OmegaDM h^2 \\simeq 0.114$.  At present, however, when $\\vrel$ is much smaller, the Sommerfeld effect becomes important, and, for example, for $m_X \\sim 2~\\tev$ and an assumed enhancement factor of $S \\sim 1000$, such annihilations are sufficient to explain the positron excesses~\\cite{Bergstrom:2009fa,Meade:2009iu}.  In Ref.~\\cite{Feng:2009hw}, we showed that in straightforward models, such large Sommerfeld enhancements cannot be self-consistently realized.  The problem is simple to state: large Sommerfeld enhancement requires strong interactions, and strongly-interacting particles annihilate too efficiently in the early Universe to be all of the dark matter.  More quantitatively, the required new force carrier interaction necessarily induces an annihilation process $XX \\to \\phi \\phi$, with cross section \\begin{equation} \\sv \\sim \\frac{\\pi \\alpha_X^2}{m_X^2} \\ . \\label{sigma} \\end{equation} Conservatively neglecting other annihilation processes, for a typical weak-scale mass $m_X$, the thermal relic density implies a typical weak coupling $\\alpha_X \\approx 0.05 \\, [m_X / 2~\\tev]$, as dictated by the WIMP miracle.  The resulting Sommerfeld enhancement now is $\\sim \\pi \\alpha_X / v\\approx 100 \\, [m_X / 2~\\tev]$, an order of magnitude too small to explain the PAMELA and Fermi signals. Alternatively, to achieve $S \\sim 1000$, one requires $\\alpha_X \\sim 1$, which implies $\\Omega_X \\sim 0.001$, two orders of magnitude smaller than the value assumed in deriving the requirement $S \\sim 1000$.  In this work, we refine our previous analysis in several ways.  In Ref.~\\cite{Feng:2009hw} we approximated the Sommerfeld enhancement $S$ by its value at $\\mphi = 0$, given by $S^0$ in \\eqref{S0}.  For massive $\\phi$, the Sommerfeld enhancement cuts off at a value proportional to $m_X/\\mphi$ and also exhibits resonant structure. Here we use numerical results (and a highly accurate analytic approximation to the numerical results~\\cite{Iengo:2009ni,Iengo:2009xf,Cassel:2009wt,Slatyer:2009vg}) for $S$.  In addition, we refine our previous work to include the effect of Sommerfeld enhancement on freeze out.  One might expect this effect to be negligible, because freeze out is typically thought to occur at $v \\sim 0.3$, when the Sommerfeld effect is insignificant.  However, annihilation continues to much later times, and Sommerfeld enhancement has an impact when the dark matter cools.  Sommerfeld effects on freeze out have been considered previously in Ref.~\\cite{Hisano:2006nn} for the case of Wino dark matter and in Refs.~\\cite{MarchRussell:2008yu,Yuan:2009bb,Dent:2009bv,Zavala:2009mi} for hidden sector dark matter.  In the analyses published after the PAMELA and Fermi excesses were reported~\\cite{Yuan:2009bb,Dent:2009bv,Zavala:2009mi}, $\\sv$ was taken as a free parameter, and the fact that it depends on the same parameters that determine $S$ was not used.  Here we make essential use of the observation that there is a irreducible contribution to $\\sv$ of the form of \\eqref{sigma}, and so constraints on the relic density bound the fundamental parameters that also determine $S$.  As we will see, the combination of these two refinements highlights a number of effects that may typically be ignored, but now require exploration.  For example, the relic density is in principle affected by the cutoff of resonant Sommerfeld enhancement, the production and decay rates of force carrier particles, the temperature of kinetic decoupling, and the efficiency of self-scattering for preserving the dark matter's thermal velocity distribution.  Very close to resonances, we also find the intriguing possibility of chemical {\\em recoupling}, in which dark matter freezes out and then melts back in, with annihilations becoming important again at late times.  This implies that exact resonances suppress, rather than enhance, indirect signals.  We examine the quantitative impact of all of these effects, and present results for the maximal Sommerfeld enhancement of indirect signals.  In the most optimistic case, that is, tuning all parameters to maximize the signal, we find that, for $\\mphi = 250~\\mev$ and $m_X = 0.1$, 0.3, and 1 TeV, the largest possible effective Sommerfeld enhancements are $\\Seff = 7$, 30, and 90, respectively. This refined analysis therefore strengthens our previous results: Sommerfeld enhancements may be significant, but they fall short of explaining the current PAMELA and Fermi cosmic ray excesses.  We then discuss various astrophysical effects that may reduce the discrepancy between the maximal enhancements derived here and those required to explain the data, and critically examine several non-minimal models and the issues that must be addressed to determine if more complicated particle physics models may provide viable explanations.  This paper is organized as follows.  In \\secref{S}, we describe our underlying model assumptions and our treatment of Sommerfeld enhancements with resonances.  In \\secref{freezeout} we analyze the effect of Sommerfeld enhancement on freeze out.  In \\secref{results} we assemble these pieces and present the results for the maximal Sommerfeld enhancement achievable now.  In \\secref{comparison}, we compare these to the current data, and discuss non-standard astrophysical effects and non-minimal particle physics models.  We present our conclusions in \\secref{conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In this work, we have investigated scenarios with Sommerfeld-enhanced annihilation of thermal dark matter and determined the largest possible indirect signals, subject to the most basic of constraints, namely, that the thermal dark matter have the correct thermal relic density.  In addition to their potential impact on cosmic $e^{\\pm}$ spectra, such scenarios may also predict observable effects in the cosmic microwave background and the diffuse extragalactic $\\gamma$-ray background, excess $\\gamma$-ray, neutrino and radio wave emission from the galactic center, and excess $\\gamma$-rays from dwarf spheroidal satellites of the Milky Way and extended regions centered on the galactic center~\\cite{Kamionkowski:2008gj,Belikov:2009qx,% Galli:2009zc,Profumo:2009uf,Slatyer:2009yq,Huetsi:2009ex,% Meade:2009iu,Cirelli:2009bb,Kanzaki:2009hf,Crocker:2010gy,% Abazajian:2010sq,Calvez:2010jq,PalomaresRuiz:2010pn}. The absence of each of these potential signatures provides significant constraints. The bounds derived here are complementary to those, in that these bounds are derived by studying not the annihilation products, but the production.  These results are therefore model-independent and robust, in that they assume only that the dark matter is produced by thermal freeze out, which is the central assumption motivating the consideration of Sommerfeld enhancements in the first place.  This analysis included the possibility of resonant Sommerfeld enhancement and the effect of Sommerfeld enhancement in the thermal relic density calculation.  The possibility of Sommerfeld enhancement at low velocities, with annihilation cross sections proportional to $\\vrel^{-1}$, or even $\\vrel^{-2}$ on resonance, implies that the thermal relic density is sensitive to the dark matter density's evolution long after conventional freeze out at $T_f \\sim m_X / 25$. As a result, the relic density depends on many aspects absent in conventional scenarios.  These include the cutoff of resonant Sommerfeld annihilation, the equilibration of force carrier particles, the temperature of kinetic decoupling $\\Tkd$, and the efficiency of self-scattering to thermalize the dark matter velocity distribution after $\\Tkd$.  These dependences may be very important.  A non-generic but intriguing example is that, on resonance, dark matter freezes out but may then later melt in, chemically recoupling for a second era of efficient annihilation.  As a result, exact resonances suppress rather than enhance indirect signals, in contrast to naive expectations.  More generally, the Sommerfeld enhancement of freeze out implies that smaller tree-level annihilation cross sections are required to give the correct relic density, and so reduces the effective Sommerfeld enhancement of indirect signals.  We have found reduction factors of a few, with several $\\sim 10 - 30\\%$ variations possible, depending on the value of $\\Tkd$ and other freeze out parameters.  For the minimal scenario analyzed here, and adopting the most optimistic assumptions to maximize indirect signals, the largest possible values of $\\Seff$ are 7, 30, and 90 for $\\mphi = 250~\\mev$ and $m_X = 0.1$, 0.3 and 1 TeV, respectively.  As shown in \\figsref{fit}{fit1gev}, such enhancements fall short of explaining the PAMELA and Fermi cosmic $e^{\\pm}$ excesses.  For the best fit point, as discussed in \\secref{comparison}, this discrepancy is over an order of magnitude and cannot be eliminated by appeals to enhanced local dark matter density or boosts from small scale structure.  Bound state effects are also unable to eliminate the inconsistency, as they are typically insignificant in the parameter regions that may potentially explain the cosmic ray anomalies, and are in any case stringently constrained by the CMB.  Along with the astrophysical uncertainties discussed in \\secref{comparison} and the complementary constraints from the cosmic microwave background and other sources, these results imply that the dark matter motivations for Sommerfeld enhancements and GeV dark forces at present are, at best, strained.  We have discussed several possibly interesting directions for further study in \\secref{comparison}. Non-minimal models will generically include additional annihilation channels, such as those involving hidden Higgs bosons.  These will generically strengthen the bounds on $\\Seff$.  We noted also that the decay of the hidden sector Higgs boson to standard model particles could be constrained by demanding consistency with the measured light element abundances produced during BBN.  Additional features and constraints particular to specific models have been noted in \\secref{comparison}.  We have focused on the $4\\mu$ final state here, which generically give the best fits to the data. Different annihilation channels, such as $4e$ final states, are optimized at lower $\\Seff$ and $m_X$, but also yield worse fits to the data according to present estimates in the literature. However, this conclusion may depend significantly on the cosmic ray propagation model parameters used. At present, no Sommerfeld models have been shown to explain the PAMELA and Fermi excesses consistent with the constraint of thermal freeze out, and we stress that complete and dedicated studies of freeze out are required to judge properly whether more complicated models can self-consistently explain PAMELA and Fermi data.  The maximal Sommerfeld enhancements derived here, although not sufficient to explain current PAMELA and Fermi excesses, are nevertheless significant enhancements.  These will be probed as PAMELA and Fermi continue to gather data, and, in the near future, at AMS on the International Space Station.  The interpretation of future data as dark matter may, of course, continue to be clouded by confusion from astrophysical uncertainties, but the diverse set of particle physics and astroparticle physics experiments that may probe weak-scale dark matter in the near future at least provides some reason for optimism."
    },
    "1005/1005.3744_arXiv.txt": {
        "abstract": "We test the effect of assumptions about stellar motion on the behavior of gravitational instabilities in protoplanetary disks around solar-type stars by performing two simulations that are identical in all respects except the treatment of the star. In one simulation, the star is assumed to remain fixed at the center of the inertial reference frame. In the other, stellar motion is handled properly by including an indirect potential in the hydrodynamic equations to model the star's reference frame as one which is accelerated by star/disk interactions. The disks in both simulations orbit a solar mass star, initially extend from 2.3 to 40 AU with a $\\varpi^{-1/2}$ surface density profile, and have a total mass of 0.14 M$_{\\odot}$. The $\\gamma = 5/3$ ideal gas is assumed to cool everywhere with a constant cooling time of two outer rotation periods.  The overall behavior of the disk evolution is similar, except for weakening in various measures of GI activity by about at most tens of percent for the indirect potential case. Overall conclusions about disk evolution in earlier papers by our group, where the star was always assumed to be fixed in an inertial frame, remain valid. There is no evidence for independent one-armed instabilities, like SLING, in either simulation. On the other hand, the stellar motion about the system center of mass (COM) in the simulation with the indirect potential is substantial, up to 0.25 AU during the burst phase, as GIs initiate, and averaging about 0.9 AU during the asymptotic phase, when the GIs reach an overall balance of heating and cooling. These motions appear to be a stellar response to nonlinear interactions between discrete global spiral modes in both the burst and asymptotic phases of the evolution, and the star's orbital motion about the COM reflects the orbit periods of disk material near the corotation radii of the dominant spiral waves. This motion is, in principle, large enough to be observable and could be confused with stellar wobble due to the presence of one or more super-Jupiter mass protoplanets orbiting at 10's AU. We discuss why the excursions in our simulation are so much larger than those seen in simulations by \\citet{rice2003a}. ",
        "introduction": "\\label{sec:intro} As reviewed in \\citet{durisen2007}, gravitational instabilities (GIs) will lead to the dynamic growth of spiral waves in protoplanetary disks whenever the Toomre parameter $Q = c_s\\kappa/\\pi G\\Sigma$ becomes less than about 1.5 to 1.7 or so. Here, $c_s$ is the sound speed, $\\kappa$ is the epicyclic frequency ($\\approx \\Omega$, the disk orbital angular frequency, in a nearly Keplerian disk), and $\\Sigma$ is the  surface mass density. As the spirals grow to nonlinear amplitude, either they will saturate at a point where radiative cooling balances the heat generated by the waves, or they will fragment into bound clumps if the radiative cooling time $t_{cool}$ is comparable to or shorter than the disk rotation period $P_{rot} = 2\\pi/\\Omega$ \\citep[e.g.,][]{gammie2001, rice2003b, rice2005, mejia2005, clarke2007}. In nonfragmenting disks, mass transport rates for realistic conditions can result in effective \\citet{shakura1973} $\\alpha$-viscosity parameters of order 10$^{-2}$ \\citep[e.g.,][]{lodato2004, boley2006}, sufficient to dominate major phases of disk evolution \\citep[e.g.,][]{vorobyov2007}. Fragmentation of disks due to fast cooling may be an important planet formation mechanism under some conditions \\citep[e.g.,][]{boss1997, boss2000, boss2001, mayer2002, mayer2004, stamatellos2009, boley2009, clarke2009, rafikov2009}.  Early simulations of GIs in disks by the Indiana Univeristy hydrodynamics group (IUHG) included a hydrodynamically active star \\citep{pickett1998, pickett2000}, but the disks we could model with active stars were only ten times larger than the stellar radius. To produce more realistically sized protoplanetary disks, we introduced a central hole and fixed the star at the center of our cylindrical computational grid \\citep{pickett2003, mejia2005, boley2006, boley2007, caiphd2006, cai2008,boley2008}. The hydrodynamics equations remained written as if the grid reference frame were inertial. With an artificially fixed central stellar potential in the inertial frame, star/disk interactions that might shift the star off the system center of mass (COM) cannot be properly modeled. Acceleration of the star and subsequent feedback into the spiral structure is explicitly suppressed. In a sense, the star has infinite inertia and becomes a sink for momentum. Under these conditions, SLING amplification \\citep{adams1989,shu1990} of spiral structure cannot be properly computed. The concern about our earlier calculations is then two-fold: 1) The physics in our simulations of GI-active disks may be inaccurate, e.g., leading to erroneous estimates of the effective $\\alpha$s for mass transport, and may lack important features, e.g., stronger and more coherent one-armed spirals. 2.) The stellar motions that we suppress could be a significant and observable signature of GI activity in disks \\citep{rice2003a}.  Authors have approached the problem of stellar motion for disk simulations in different ways. In Smoothed Particle Hydrodynamics (SPH) simulations of GIs \\citep{rice2003b,mayer2004,lodato2004,stamatellos2008,forgan2009}, the stellar motion is included automatically by treating the star as a central sink particle that is smoothed differently from the rest of the SPH particles. However, these simulations, as in \\citet{rice2003a}, can  exhibit a large initial accretion rate onto the central object. How the transfer of angular momentum from the disk material to the central object is handled could be quite important to the star's motion. The only global 3D grid-based hydrodynamics scheme other than the IUHG code so far used for published simulations of GIs in protoplanetary disks around solar-type stars is the Eulerian spherical-grid code of Alan Boss. Although he does not explicitly integrate the stellar equation of motion, \\citet{boss2000} allows the central protostar to wobble in response to the growth of nonaxisymmetry in the disk by repositioning the star to keep the system COM at the grid center. There is no guarantee with this scheme that the star's displacement is a true response to Newtonian reaction forces applied by the disk. The star's motion could simply represent an accumulation of numerical error in the location of the COM. In principle, then, this treatment is no better than what was done in the IUHG simulations. In order to explore the effect of the star/disk interaction, we have now implemented the indirect potential method \\citep[e.g.,][]{nelson2000b}. This effectively puts our simulations into the accelerated reference frame of the star and therefore properly accounts for stellar motion while keeping the star at the center of our computational grid.  In this article, we present the results of a simulation for a nonfragmenting disk using the indirect potential to treat stellar motion and compare it to results of an identical simulation with an artificially fixed central star from Mej\\'ia et al. (2005). In \\S\\ref{sec:results}, we first present an overall qualitative comparison, followed by a detailed analysis of the stellar motion and of differences in the disk behaviors. We compare our results to those obtained via other numerical methods and briefly discuss possible observational consequences in \\S\\ref{sec:compare}. Finally, \\S\\ref{sec:conclude} presents our main conclusions.  \\begin{figure*} \\includegraphics{fig1.eps} \\caption{Midplane densities in logarithmic scale of the indirect (top) and fixed (bottom) simulations. The axes have units of AU and the time in the upper right is given in ORPs. The small white x in the indirect panels is the position of the system center of mass (disk plus star). The left panels show the saturation of the initial growth in the burst phase. The right panels are representative of the asymptotic behavior.} \\label{SMfinalcomp} \\end{figure*} ",
        "conclusions": "\\label{sec:conclude}  The two major goals of this paper were: 1) to determine how much a proper treatment of stellar motion affects qualitative and quantitative conclusions made in previous studies by our group and 2) to decide whether stellar motion due to star/disk interaction could lead to observable consequences. We find that, although the simulations do show a weakening of disk self-torques and nonaxisymmetric structure by at most about tens of percent, simulations with and without proper treatment stellar motion are overall very similar. They exhibit the same evolutionary phases, and they both settle into similar quasi-steady asymptotic states where cooling is balanced by GI heating. The main qualitative difference is a shift of dominance from an even mode (four-armed) to an odd mode (five-armed) during the burst phase when the star's motion is properly treated. This would only matter in a case where the initial conditions at the time of onset for GIs were well understood for a particular application. We conclude that IUHG studies using a fixed star \\citep[e.g.,][]{pickett2003,mejia2005,cai2006,boley2006,boley2007,boley2008,cai2008} have produced valid and reliable results regarding the overall behavior and evolution of GI-active disks.  When we account for stellar motion through an indirect potential, the motion of the star appears to be due to the nonlinear interaction of discrete odd and even global spiral modes, both during the burst and in the asymptotic phase. These interactions involve substantial disk mass at radii of about 30 AU (burst) and 24 AU (asymptotic phase) with the result that both astrometric and radial velocity disturbances are fairly large. For the asymptotic phase, we estimate about $16 ~\\mathrm{\\mu as/yr}$ (at 100 pc) and 30 m/s, respectively. The time signatures can involve multiple periodicities, but it would be difficult to distinguish this from a stellar response to the presence of multiple massive planets or protoplanets. Our results are in contrast to the much more modest stellar motions reported by \\citet{rice2003a}, under conditions where the cooling times decrease inward in the disk, there is no burst, GIs are not dominated by a few discrete global modes, and there is significant accretion from the inner disk. This means that stellar motion induced by GIs is sensitive to various aspects of disk structure and physics. Further study is clearly warranted, especially if outbursts of GIs are common during the early phases of disk formation when disks are large and massive \\citep[e.g.][]{laughlin1994,vorobyov2006,stamatellos2009,boley2009}."
    },
    "1005/1005.3602_arXiv.txt": {
        "abstract": "{The CoRoT satellite has recently discovered the transits of a telluric planet across the disc of a late-type magnetically active star dubbed CoRoT-7, while a second planet has been detected {{ after filtering out the radial velocity (hereafter RV) variations due to stellar activity.}} } { {{ We investigate the magnetic activity of CoRoT-7 and use the results for a better understanding of the impact of magnetic activity on stellar RV variations. }} } {{{ We derive the longitudinal distribution of active regions on CoRoT-7 from a maximum entropy spot model of the CoRoT light curve. Assuming that each active region consists of dark spots and bright faculae in a fixed  proportion, we synthesize the expected RV variations. }} } {Active regions are mainly located at three active longitudes which {{ appear to }} migrate at different rates, probably as a consequence of surface differential rotation, {{ for which a lower limit of $\\Delta \\Omega / \\Omega = 0.058 \\pm 0.017$ is found}}. The synthesized activity-induced RV variations  reproduce the  amplitude of the observed RV curve and are used to study the impact of stellar activity on planetary detection.  } {{ {{In spite of the non-simultaneous CoRoT and HARPS observations, our study confirms the validity of the method previously adopted to filter out RV variations induced by stellar activity. We find a false-alarm probability $ < 10^{-4}$ that the RV oscillations attributed to CoRoT-7b and CoRoT-7c are spurious effects of noise and activity. Additionally, our model suggests that other periodicities found in the observed RV curve of CoRoT-7 could be explained by active regions  whose visibility is modulated by a differential stellar rotation with periods ranging from 23.6 to 27.6 days.}}  }} ",
        "introduction": "CoRoT (Convection, Rotation and Transits) is a photometric space experiment devoted to asteroseismology and the search for extrasolar planets by the method of transits \\citep{Baglinetal06}. It has recently discovered two telluric planets around a late-type star called CoRoT-7, the innermost of which transits across the disc of the star and has an orbital period of only 0.8536 days, while the other has a period of 3.69 days and does not show transits \\citep{Legeretal09,Quelozetal09}. The mass of the inner planet CoRoT-7b is $4.8 \\pm 0.8$ M$_{\\oplus}$, while  that of CoRoT-7c is $8.4 \\pm 0.9$ M$_{\\oplus}$, assuming they are on coplanar orbits \\citep{Quelozetal09}. The relative depth of the transit of CoRoT-7b is $\\Delta F/F = (3.35 \\pm 0.12) \\times 10^{-4}$ in the CoRoT white passband (see Sect.~\\ref{observations}), leading to a planetary radius of $1.68 \\pm 0.09$ R$_{\\oplus}$.  CoRoT-7 is an active star showing a rotational modulation of its optical flux with an amplitude up to $\\sim 0.03$ mag {and} a period of $\\sim 23.5$ days \\citep{Quelozetal09}. The discovery of CoRoT-7c and the measurement of the mass of CoRoT-7b have become possible with the HARPS  spectrograph after filtering out the apparent variations of the stellar RV induced by magnetic activity which have an amplitude of $\\approx 40$ m s$^{-1}$, i.e., $5-6$ times greater than the wobbles produced by the gravitational pull of the planets. {Therefore, a more detailed model of the stellar activity is needed to improve RV measurements and to obtain better planetary parameters. This would also rule out the possibility that a quasi-periodic RV signal is due to rotating spots on the stellar surface instead of a planetary companion as, e.g., in the case of HD~166435 \\citep{Quelozetal01}. }  We apply the same spot modelling approach used for the high-precision light curves of CoRoT-2 \\citep{Lanzaetal09a} and CoRoT-4 \\citep{Lanzaetal09b}, deriving the longitudinal distribution of photospheric brightness inhomogeneities and their evolution {during} the $\\sim 130$ days of CoRoT observations. The different rotation rates of active regions  allow us to estimate a lower limit for the amplitude of  the differential rotation of CoRoT-7. From our spot maps we  synthesize the activity-induced RV perturbations using a simple model for the distortion of the line profiles, as described in Sect.~\\ref{line_prof_model}. A Fourier analysis of such a synthetic time series  {yields typical periods} induced by stellar activity, allowing us to discuss its impact on the detection of planets around CoRoT-7. ",
        "conclusions": "\\label{conclusions}  We have applied the spot modelling method introduced by \\citet{Lanzaetal09a,Lanzaetal09b} to the {lightcurve of the } planet-hosting star CoRoT-7. Adopting a facular-to-spotted area ratio $Q=1$, we have found evidence of three persistent active longitudes within which individual active regions form and decay with lifetimes ranging from a couple of weeks to a few months. The  active longitudes migrate in the adopted reference frame at different rates which suggests  a surface differential rotation with a relative amplitude of at least $0.058 \\pm 0.017$. {The overall spotted area  reaches a maximum of 1.6 percent of the whole surface if the contrasts of dark spots and faculae are adopted as solar-like. } The residuals of the best fit  indicate a population of  small spots with areas of the order of $5\\times 10^{-4}$ of the whole stellar surface that evolve on a typical timescale of $\\sim 4-5$ days. Moreover, we find  evidence of an oscillation of the optical flux {with} a semiamplitude $\\sim 6 \\times 10^{-5}$ mag with the orbital period of the planet CoRoT-7c.  We introduce a  model to simulate the apparent RV variations induced by magnetic activity which includes the effect of surface brightness inhomogeneities as well as a reduction of the macroturbulence and convective blueshifts in  active regions. We synthesize a time series for the RV variations of CoRoT-7 adopting the spot model obtained from CoRoT white band light curve and assuming that active regions have solar-like contrasts and are located within $\\pm 20^{\\circ}$ from the equator. Unfortunately, HARPS and CoRoT observations are not simultaneous, so any model based on the latter does not constrain, rigorously speaking, the observed RV  variations.  Nevertheless, {an extrapolation of our model matches the observed RV amplitude.}  We fit the RV  time series with a linear combination of three sinusoids at the rotation frequency and its first two harmonics, following  \\citet[][ Sect.~4]{Quelozetal09}, and confirm that their method is appropriate to filter out most of the activity-induced RV  variations. The residual amplitudes at the orbital periods of the planets CoRoT-7b and CoRoT-7c are  used to estimate false-alarm probabilities that the RV oscillations attributed to the planets are spurious effects induced by stellar activity. For both planets we find a false-alarm probability lower than $10^{-4}$.  The possible presence of a third planet with a period of 9.03 days cannot be confirmed by means of our approach because this period could be associated to an harmonic of the signal with a period of 27.6 days which can be produced by spots at high latitude on a differentially rotating star. We shall consider the effects of surface differential rotation in future models of the RV variations to investigate this issue."
    },
    "1005/1005.1607_arXiv.txt": {
        "abstract": "We study the contribution of massive dominantly sterile neutrinos, $N$,  to the Lepton Number and Lepton Flavor Violating (LNV, LFV) semileptonic decays of $\\tau$ and  $B, D, K$-mesons. We focus on special domains of sterile neutrino masses $m_{N}$ where it is close to its mass-shell. This leads to an enormous resonant enhancement of the decay rates of these processes. This allows us to derive stringent limits on the sterile neutrino mass $m_{N}$ and its mixing matrix elements $|U_{e N}|, |U_{\\mu N}|, |U_{eN} U_{\\tau N}|, |U_{\\tau N} U_{\\tau N}|$, by  carrying out a joint analysis of the existing experimental bounds on the decay rates of the studied processes. In contrast to other approaches in the literature, our extraction of these limits is free from ad hoc assumptions on the relative size of the sterile neutrino mixing parameters.  Special attention was paid to the limits on meson decays with $e^{\\pm}e^{\\pm}$ in final state, derived from non-observation of $0\\nu\\beta\\beta$-decay.  We point out that observation of these decays may, in particular,  shed light on CP violation in the neutrino sector. ",
        "introduction": "\\label{sec-1}   Neutrino oscillation experiments have shown conclusively that neutrinos are massive, although very light, particles mixing with each other. Moreover, neutrino oscillations is the first and so far  the only observed phenomenon of lepton flavor violation (LFV). Theoretically LFV is only possible if neutrinos are massive. Being electrically neutral neutrinos can be either Dirac or Majorana type particles. Majorana neutrino masses violate total lepton number conservation and, therefore, can induce lepton number violating (LNV) processes. Up to date such processes have not yet been observed experimentally.  Searching for LNV processes is a challenging quest, pointed to probe the nature of neutrinos and answer the question about wether they are Majorana or Dirac particles.  However it  is well known that in Standard Model (SM) extensions with very light and very heavy neutrinos, the rates of both LFV and LNV processes are  so small that their experimental observation turns out to be unrealistic. Perhaps the only lucky exception is neutrinoless double beta decay ($0\\nu\\beta\\beta$). The experiments searching for this LNV  process are believed to reach even such small neutrino masses \\cite{dbd-probe, KlapdorKleingrothaus:2000sn} as those which are relevant for neutrino oscillations, and to probe the heavy Majorana neutrino sector up to the masses of several TeV \\cite{Bamert:1995,Benes:2005hn}.   On the other hand there are other extensions of the SM in which both LFV and LNV effects could be significant for various reasons. In particular there may contain new massive particles like moderately heavy Majorana neutrinos, as in the low-scale versions of the see-saw scenario, or other new particles such as neutralinos, heavy Majorana particles mixing with neutrinos in the framework of SUSY models with R-parity violation.  These new heavy particles may lead to observable effects beyond the light neutrino sector, since they could manifest themselves indirectly via their virtual contribution to processes involving ordinary particles. If  these new particles are not very heavy lying within the mass reach of the running of forthcoming experiments, they can be searched for directly among the products of colliding particles. Also they may contribute to some decays as virtual, but nearly on-mass-shell states, leading to huge resonant enhancement of the corresponding decay rates.  Here we consider an extended see-saw scenario, including $n$ species of SM singlet right-handed neutrinos $\\nu^{\\,\\prime}_{Rj}=(\\nu^{\\,\\prime}_{R1},...\\nu^{\\,\\prime}_{Rn})$, besides the three left-handed weak doublet neutrinos $\\nu^{\\,\\prime}_{Li} = (\\nu^{\\,\\prime}_{Le},\\nu^{\\,\\prime}_{L\\mu},\\nu^{\\,\\prime}_{L\\tau})$. The general mass term for this set of fields can be written as \\ba - \\frac{1}{2} \\overline{\\nu^{\\,\\prime}} {\\cal M}^{(\\nu)} \\nu^{\\,\\prime c} + \\mbox{h.c.} & = & - \\frac{1}{2} (\\bar\\nu^{\\,\\prime}_{_L}, \\overline{\\nu_{_R}^{\\,\\prime c}}) \\left(\\begin{array}{cc} {\\cal M}_L & {\\cal M}_D \\\\ {\\cal M}^T_D  & {\\cal M}_R \\end{array}\\right) \\left(\\begin{array}{c} \\nu_{_L}^{\\,\\prime c} \\\\ \\nu^{\\,\\prime}_{_R}\\end{array}\\right) + \\mbox{h.c.} \\\\ &=&-\\frac{1}{2} (\\sum_{i=1}^{3} m_{\\nu_i} \\overline{\\nu^c_i}\\nu_i +\\sum_{j=1}^{n} m_{\\nu_j} \\overline{\\nu^c_j}\\nu_j  )+ \\mbox{h.c.} \\ea Here ${\\cal M}_L, {\\cal M}_R$ are $3\\times 3$ and $n\\times n$ symmetric Majorana mass matrices, and ${M}_D$ is a $3\\times n$ Dirac type matrix. Rotating the neutrino mass matrix to the diagonal form by a unitary transformation \\ba && U^T {\\cal M}^{(\\nu)}U = \\textrm{Diag}\\{m_{\\nu 1}, \\cdots ,m_{\\nu_{3+n}}\\} \\label{mass-eigen} \\ea we end up with $3+n$ Majorana neutrinos with masses $m_{v_1}, \\cdots, m_{v_{3+n}}$. In special cases there may appear among them pairs with masses degenerate in absolute values. Each of these pairs can be collected into a Dirac neutrino field. This situation corresponds to conservation of certain lepton numbers assigned to these Dirac fields.  Neutrino interactions are represented by the SM Charged (CC) and Neutral Current (NC) Lagrangian terms.  In the mass eigenstate basis they read \\ba{CC-NC} {\\cal L} = \\frac{g_2}{\\sqrt{2}}\\sum_{i}\\  U_{l i}\\ \\bar l \\gamma^{\\mu} P_L \\nu_i\\  W^-_{\\mu} +  \\frac{g_2}{2 \\cos\\theta_W}\\  \\sum_{\\alpha, i, j}U_{\\alpha j} U_{\\alpha i}^*\\  \\bar\\nu_i \\gamma^{\\mu} P_L \\nu_j\\ Z_{\\mu}, \\ea where $l = e,\\mu,\\tau$ and $i=1,..., n+3 $.   In this scenario the light neutrinos have the mass scale $\\mathcal{M}_D^2/\\mathcal{M}_R$, and as a result there should be also heavy Majorana neutrinos ($N$) with mass scale $\\mathcal{M}_R$. In particular, they may include moderately heavy neutrinos $N$ in the MeV-GeV domain, which we are going to consider in the next sections. If they exist they may contribute to some LNV and LFV processes as intermediate on-mass-shell states. This would result in enormous resonant enhancement of their contribution to these processes  and would allow one to set stringent limits on heavy neutrino masses and mixing from non-observation of the corresponding processes. The presence or absence of these neutrino states, which are conventionally called sterile neutrinos, is a subject to be revealed by experimental searches. Having these arguments in mind we study neutrino states $N$ with masses $m_{N}$ in this  region of values and examine their phenomenological impact on the LFV and LNV decays of $\\tau$ and some mesons. The phenomenology of sterile neutrinos have been studied in the literature in different contexts and from different points of view. Their resonant contributions to $\\tau$ and meson decays have been studied in Refs.  \\cite{K-decPaper}, \\cite{tau-decPaper}, \\cite{Ivanov:2004ch}. Another potential process to look for  sterile Majorana neutrinos is like-sign dilepton production in hadron collisions \\cite{Almeida:2000pz}, \\cite{Panella:2001wq}, \\cite{Han:2006ip}, \\cite{Kovalenko:2009td}.  Possible implications of sterile neutrinos have been also studied in LFV muonium decay and high-energy muon-electron scattering \\cite{Cvetic:2006yg}. (for a recent review of sterile neutrino phenomenology see \\cite{Atre:2009rg}). The present paper contributes to some still uncovered aspects of the phenomenology of sterile neutrinos.  The paper is organized as follows. In the next section \\ref{sec-2} we present decay rate formulas for $\\tau$ and meson LFV and LNV decays and their reduced forms, in the resonant domains of neutrino masses. In sec.\\ref{sec-3} we discuss theoretical uncertainties in the calculation of the total decay width of sterile neutrinos and present our ``inclusive'' approach based on Bloom-Gilman duality. This approach allows one to avoid large uncertainties related to the heavy meson decay constants, which are  typical in the conventional channel-by-channel approach. In sec.\\ref{sec-4} we discuss the extraction of the limits on  the active-sterile neutrino mixing matrix elements $|U_{e N}|, |U_{\\mu N}|, |U_{eN} U_{\\tau N}|, |U_{\\mu N} U_{\\tau N}|$ and present the corresponding exclusion plots. In sec. \\ref{sec-5} we discuss the $0\\nu\\beta\\beta$-decay constraints on sterile neutrino parameters and their possible implications for CP-violation in the neutrino sector. Here we also give predictions for the rates of some LNV and LFV decays of $K, D, D_{s},B, B_{c}$-mesons. Sec.VI summarizes our main results. ",
        "conclusions": "\\label{sec-6} We have studied LFV and LNV decays of $\\tau$ and $B, D, K$ mesons  mediated by heavy Majorana neutrinos. We focussed on the dominant mechanism via resonant Majorana neutrino production.  On the basis of joint analysis of experimental limits on certain sets of LNV and LFV decays with intersecting resonant regions we extracted upper bounds on the heavy Majorana neutrino mass $m_{N}$ and mixing $|U_{eN}|, |U_{\\mu N}|, |U_{e N} U_{\\tau N}|, |U_{\\mu N} U_{\\tau N}|$ without additional assumptions on these parameters. Our limits derived from the above mentioned decays are shown in Figs. \\ref{figure2}-\\ref{figure4}. As seen they are less stringent than the corresponding limits derived in the literature. To our best knowledge the latter typically involve various ad hoc assumptions on the relative size of the heavy neutrino mixing matrix elements. Therefore,  it is natural that our limits, not referring to any kind of such assumptions, are less stringent. The mixing $U_{\\tau N}$ is left unconstrained by the present experimental data, and requires searches for processes with $\\tau\\tau$ in the final state.  In our analysis we used the method of calculation of the total decay rate of heavy Majorana neutrino on the basis of an inclusive approach. Instead of a detailed calculation of the sum of  all the possible decay channels with different hadrons in the final state, we approximated the hadronic final states by a $q\\bar{q}$-pair,  as  suggested by Bloom-Gilman duality. This allows us to avoid large theoretical uncertainties in the decay constants of heavy mesons.   Special emphasis has been made on the stringent limits on the heavy Majorana neutrino from non-observation  of $0\\nu\\beta\\beta$-decay, from which we also derived indirect limits on various LNV and LFV processes. Our results, displayed in Fig. \\ref{figure2} and  Table II,  have been obtained assuming no CP-violation in the neutrino sector. As seen from Table II, our expectations for the prospects of observation of the decays with both $e^{\\pm}e^{\\pm}$ and $e^{\\pm}e^{\\mp}$ in the final states are pessimistic. One of the messages of the present paper is that despite of such a discouraging prediction these processes are worth searching. The point is that any observation of such decays above the limits set by the CP conserving $0\\nu\\beta\\beta$ decay constraints would most likely  point to the existence of heavy Majorana neutrinos in the resonant regions of these processes, and CP-violation in the neutrino sector, as explained in sec. V.  We applied our limits on heavy Majorana neutrino mass $m_{N}$ and its mixing $U_{lN}$ for a prediction of upper limits on the  rates of as yet experimentally unconstrained LFV and LNV decays of $B$ and $D$ mesons. These limits are shown in Table II, and indicate that only the experiments searching for the LNV and LFV decays of $D$ and $D_{s}$ are rather close to probe the rates of these processes not excluded by our analysis.  Searches for other decays are far away from this perspective. \\\\[5mm] {\\bf Acknowledgements} This work is supported by FONDECYT (Chile) under projects 1100582, 1100287 and Centro-Cient\\'\\i fico-Tecnol\\'{o}gico de Valpara\\'\\i so PBCT ACT-028."
    },
    "1005/1005.2481_arXiv.txt": {
        "abstract": "s{Probing Gaussianity represents one of the key questions in modern cosmology, because it allows to discriminate between different models of inflation. We test for large-scale non-Gaussianities in the cosmic microwave background (CMB) in a model-independent way. To this end, so-called first and second order surrogates are generated by first shuffling the Fourier phases belonging to the scales not of interest and then shuffling the remaining phases for the length scales under study. Using scaling indices as test statistics we find highly significant signatures for both non-Gaussianities and asymmetries on large scales for the WMAP data of the CMB. We find remarkably similar results when analyzing different ILC-maps based on the WMAP five and seven year data. Such features being independent from the map-making procedure would disfavor the fundamental principle of isotropy as well as canonical single-field slow-roll inflation - unless there is some undiscovered systematic error in the collection or reduction of the CMB data or yet unknown foreground contributions.} ",
        "introduction": "One of the key questions in cosmology is to probe the Gaussianity of the primordial density fluctuations, because it allows to discriminate between different models of inflation. While the simplest model of inflation, namely single-field slow-roll inflation\\cite{Guth81}, predicts that the temperature fluctuations of the cosmic microwave background (CMB) correspond to a (nearly) Gaussian, homogeneous, and isotropic random field, more complex models may give rise to non-Gaussianity (NG). Models in which the Lagrangian is a general function of the inflaton and powers of its first derivative can lead to scale-dependent non-Gaussianities, if the sound speed varies during inflation\\cite{Armendariz99,Garriga99}. Similarly, string theory models that give rise to large non-Gaussianity have a natural scale dependence \\cite{Loverde08a}. Possible deviations from Gaussianity  have been investigated in studies based on e.g. the WMAP data of the CMB (see e.g. Komatsu et al.\\cite{Komatsu09a} and references therein) and claims for the detection of non-Gaussianities and other anomalies, like hemispherical power asymmetry, the 'axis of evil' , the Cold Spot etc.  have been made. However, most of the tests on non-Gaussianities do not take into account any possible scale-dependency. In this contribution we apply the formalism of surrogate data sets, which was recently  adapted to CMB data analysis \\cite{Raeth09}, to test for non-Gaussianities on large scales in the WMAP data. Special emphasis is put on a comparison of the WMAP five year data with the WMAP seven year data. ",
        "conclusions": "In conclusion, we detect highly significant signatures for asymmetries and non-Gaussianities for large scales ($l<20$) in the WMAP five and seven year data. The increase of the signal with decreasing noise as well as the very high correlations between the significance maps points towards an intrinisic nature of the detected anomalies, which are independent of the map making procedure. Such features would disfavor the fundamental principle of isotropy as well as canonical single-field slow-roll inflation - unless there is some undiscovered systematic error in the collection or reduction of the CMB data or yet unknown foreground contributions. Thus, further tests are required to rule out other systematic effects as origin of the detected anomalies."
    },
    "1005/1005.4058_arXiv.txt": {
        "abstract": "We study the growth of dark matter halos in the concordance \\LCDM{ }cosmology using several N-body simulations of large cosmological volumes. We construct merger trees from the Millennium and Millennium-II simulations, covering the ranges $10^{9}-10^{15}\\Msun$ in halo mass and $1-10^{5}$ in merger mass ratio. Our algorithm takes special care of halo fragmentation and ensures that the mass contribution of each merger to halo growth is only counted once. This way the integrated merger rate converges and we can consistently determine the contribution of mergers of different mass ratios to halo growth. We find that all resolved mergers, up to mass ratios of $10^5:1$, contribute only $\\approx60\\%$ of the total halo mass growth, while major mergers are subdominant, e.g.~mergers with mass ratios smaller than $3:1$ ($10:1$) contribute only $\\approx20\\%$ ($\\approx30\\%$). This is verified with an analysis of two additional simulation boxes, where we follow all particles individually throughout cosmic time. Our results are also robust against using several halo definitions. Under the {\\it assumption} that the power-law behaviour of the merger rate at large mass ratios can be extrapolated to arbitrarily large mass ratios, it is found that, independent of halo mass, $\\approx40\\%$ of the mass in halos comes from genuinely smooth accretion of dark matter that was never bound in smaller halos. We discuss possible implications of our findings for galaxy formation. One implication, assuming as is standard that the pristine intergalactic medium is heated and photoionized by UV photons, is that all halos accrete $>40\\%$ of their baryons in smooth \"cold\" $T\\gtrsim10^4K$ gas, rather than as warm, enriched or clumpy gas or as stars. ",
        "introduction": "\\label{s:intro} The way galaxies gain their mass affects almost every aspect of galaxy evolution. The distinction between, e.g., accretion of gas versus stars, spherical versus filamentary accretion, or clumpy versus smooth accretion will result in very different star-formation histories, colours, morphologies, angular momentum contents and sizes. Mergers are believed to play an important role in the evolution of galaxies, in particular of elliptical galaxies via the morphological transformation from disk-dominated galaxies to spheroids (e.g.~\\citealp{ToomreA_77a,BarnesJ_92a,NaabT_03a,NaabT_06a,CoxTJ_06a,HopkinsP_07b,NaabT_09a,ConseliceC_03a,BellE_06b,LotzJ_08a}). Both theoretical and observational work have also emphasised the importance of smooth accretion of gas, in particular for the buildup of massive galaxies at high redshift and for the subsequent evolution of disk galaxies (e.g.~\\citealp{WhiteS_91a,MuraliC_02a,KeresD_05a,OcvirkP_08a,DekelA_09b,GoerdtT_10a,DaddiE_07a,FoersterSchreiberN_09a,ConseliceC_09a,KauffmannG_10a}). The gas can be accreted in a spherically symmetric mode of cooling halo gas or in a filamentary mode directly from the cosmic web.  In the cold dark matter paradigm for structure formation \\citep{BlumenthalG_84a}, galaxy formation is closely related to the formation of dark matter halos \\citep{WhiteS_78a}, albeit in a complex way. Baryons follow the flow of the gravitationally-dominating dark matter and fall into dark matter halos, and galaxy mergers follow the mergers of their host dark matter halos. Understanding the process of halo mass assembly is an important step towards a better understanding of galaxy formation. In this paper we use high resolution dark matter simulations to study the relative importance of mergers versus smooth accretion for the buildup of halos and galaxies.  The non-linear nature of the evolution of gravitational instabilities into virialised structures makes N-body simulations the most reliable tool available for studying the mass buildup of dark matter halos. Massive dark matter particles in N-body simulations are believed to give a good representation of the coarse-grained phase-space structure of real dark matter. However, the interpretation of dark matter simulations is subject to some freedom and uncertainty in subsequent steps in making the connection to galaxies, such as the definition of a dark matter halo (e.g.~\\citealp{WhiteM_01a,WhiteM_02a,CohnJ_07a,TinkerJ_08a}) and algorithms of merger tree construction.  One of the largest simulations used so far for studying the growth of dark matter halos is the Millennium Simulation (\\citealp{SpringelV_05a}; hereafter MS). \\citet{FakhouriO_07a} analysed the MS and found that the dark matter halo merger rate has a nearly universal form that can be separated into its dependencies on mass ratio, descendant mass, and redshift. They presented three algorithms for merger tree construction that result in merger rates differing by $\\approx25\\%$. In \\citet{GenelS_09a} (hereafter G09) we introduced a novel merger tree construction algorithm we termed \"splitting\" (see also \\citealp{FakhouriO_08b} and \\citealp{MallerA_06a}) that incorporates the complicated process of halo fragmentation and re-merging such that a given pair of halos is never considered to merge more than once. This resulted in a different set of parameters for the \\citet{FakhouriO_07a} global fitting formula for the merger rate extracted from our trees of the MS.  Thanks to ever increasing computation power, the dynamic range of N-body simulations increases too. As a result, several recent works have studied the growth rate of dark matter halos. When \\citet{StewartK_08a} integrated the merger contribution and extrapolated to include {\\it all} unresolved mergers, they found that only $\\approx50-70\\%$ of the final mass of halos was assembled by mergers. Rather dissimilar results were found by \\citet{MadauP_08a}, who investigated the formation of one halo in a cosmological 'zoom-in' simulation, and by \\citet{AnguloR_09a}, who investigated the growth of dark matter halos using high-resolution extended Press-Schechter (EPS; \\citealp{PressW_74a,BondJ_91a,BowerR_91a}) trees. These works concluded that the mass accretion of dark matter halos is largely dominated by mergers. In this paper we perform a similar analysis to that of \\citet{StewartK_08a} with a larger dynamic range using a combination of the MS and the higher resolution Millennium-II Simulation (\\citealp{Boylan-KolchinM_09a}; hereafter MS2), and emphasise the importance of choosing an appropriate merger tree construction algorithm. We show that all resolved mergers contribute $\\lesssim60\\%$ to halo mass growth, and suggest that $\\approx40\\%$ of the accretion rate may be genuinely smooth. We also provide further support for these results by directly following the histories of dark matter particles in two smaller cosmological boxes, and make detailed comparisons with the previous results mentioned above.  This paper is organised as follows. In \\S\\ref{s:merger_trees_method} we review the Millennium simulations, structure identification and merger-tree construction. In \\S\\ref{s:merger_rate} we provide a fitting function that describes the halo merger rate, and in \\S\\ref{s:halo_growth} we discuss the contribution of mergers and smooth accretion to halo mass buildup. In \\S\\ref{s:comparison_EPS} we compare our results to the EPS model. In \\S\\ref{s:particles_method} we describe our direct analysis of dark matter particles histories in two additional simulations, and in \\S\\ref{s:particles_results} we present the results of that analysis. We discuss implications of our results to galaxy formation in \\S\\ref{s:galaxy_formation} and compare our results to previous work in \\S\\ref{s:previous_work}. We summarise in \\S\\ref{s:summary}. ",
        "conclusions": "\\label{s:summary} In this paper we calculate the merger rates of dark matter halos and we investigate the role of smooth accretion versus mergers  to their growth. We extract the merger rates and accretion histories from the Millennium and Millenium-II simulations, combined with two additional, smaller, cosmological simulations. We use the \"splitting\" merger tree construction algorithm described in \\S\\ref{s:merger_trees_method} and in G09, and verify its reliability by reproducing our results by following individual particle histories alone, independent of merger tree construction algorithms.  We find that the contributions of all resolved mergers (up to mass ratios $\\approx10^{5}$) to the total growth rate of halos do not exceed $60\\%$, regardless of halo mass and redshift. Most of the {\\it merger} contribution comes from small mass ratio (\"major\") mergers (e.g.~$1<x<10$ contribute $\\approx30\\%$ of the total growth), while \"very minor\" mergers add very little mass (e.g.~$10^{3}<x<10^{5}$ contribute just a few percent to the total halo growth). We find that the power-law index of the merger rate is such that if the merger rate is extrapolated beyond the maximum resolved mass ratio $\\approx10^{5}$, the total contribution of {\\it all} mergers saturates at $\\approx60\\%$. This suggests that a significant mass fraction of halos may be accreted in a genuinely smooth way.  Our results have important implications for galaxy formation models and the modes of baryonic accretion. If $\\approx40\\%$ of the dark matter that is accreted onto a halo was never previously bound in any merging smaller halos, then at least $\\approx40\\%$ of the baryons must also be accreted smoothly - as gas that was never previously heated by feedback processes or converted to stars. For the baryons, this $\\approx40\\%$ is a strong lower limit since halos with $T_{vir}<10^4K$ likely cannot retain their gas. The common assumption that halos with $T_{vir}<10^4K$ cannot retain their gas also makes our results insensitive to the limited resolution of the simulations we use, because we resolve all halos above this limit (at $z\\lesssim3$). The implication is that a very large fraction of the baryonic matter falling into a halo must be pristine \"cold\" IGM gas, with $T\\gtrsim10^4K$ set by IGM photoheating. This gas is not expected to have formed stars or to have become significantly enriched with metals, no matter what the star-formation efficiency and history of the baryons in any of the merging halos."
    },
    "1005/1005.1161_arXiv.txt": {
        "abstract": "We report the discovery of Quasi Periodic Oscillations (QPO) at 0.02 Hz in a transient high mass X-ray binary pulsar KS 1947+300 using {\\em RXTE}-PCA. The QPOs were detected during May-June 2001, at the end of a long outburst. This is the 9th transient accretion powered high magnetic field X-ray pulsar in which QPOs have been detected and the QPO frequency of this source is lowest in this class of sources. The unusual feature of this source is that though the outburst lasted for more than 100 days, the QPOs were detected only during the last few days of the outburst when the X-ray intensity had decayed to 1.6\\% of the peak intensity. The rms value of the QPO is large, $\\sim15.4\\pm1.0\\%$  with a slight positive correlation with energy. The detection of QPOs and strong pulsations at a low luminosity level suggests that the magnetic field strength of the neutron star is not as high as was predicted earlier on the basis of a correlation between the spin-up torque and the X-ray luminosity. ",
        "introduction": "KS 1947+300 is an accretion powered X-ray pulsar which was first detected in June 1989 in observations made with the TTM coded-mask imaging spectrometer aboard the {\\em Kvant} module of the {\\em Mir Space Station} (Skinner 1989). The 2-27 keV flux was $70\\pm10$ mCrab. In two months the flux from the source faded by a factor of seven. Its X-ray spectrum could be described by a power law with photon index $\\Gamma = 1.72\\pm0.31$ and a hydrogen column density $N_{H} = (3.4\\pm3.0) \\times 10^{22}$ (Borozdin et al. 1990). The coordinates of the source were determined to be : R.A.$ = 19^{h}47^{m}35^{s}.2$, Dec.= $+30^{0} 04'47''$ (Eq. 1950.0). In April 1994, the Burst and {\\em Burst and Transient Source Experiment} (BATSE) aboard the {\\em Compton Gamma Ray Observatory (CGRO)} detected 18.7 s pulsations from an X-ray source within a few degrees of KS 1947+300. The newly detected source GRO J1948+32 was later found to be same as KS 1947+300 (Swank and Morgan 2000). The optical counterpart is a V=$14.2$ B0 Ve star with  moderate reddening that indicates the distance to the system to be about 10 kpc (Negueruela et al. 2003). \\\\  One large outburst and several smaller outbursts of KS 1947+300 have been observed with the {\\em Rossi X-ray Timing Explorer (RXTE)}-ASM. The first outburst of this source was detected by {\\em RXTE}-ASM in October 2000 (Levine and Corbet 2000). Following this the intensity declined, but the source became highly active again in November 2000. The outburst reached its peak in February 2001 and slowly declined till June 2001. Based on data acquired with {\\em RXTE-PCA} during the 2000-2001 outburst, Galloway et al.  (2004) determined the orbital parameters of the binary: the orbital period $P_{orb}=40.415\\pm0.010$ d, the projected semi major axis $a_{x}sini = 137\\pm3$ lt-sec and eccentricity $e = 0.033\\pm0.013$. Glitches are mainly observed in Anomalous X-ray Pulsars (AXP) and radio pulsars. But the {\\em RXTE} analysis of this source revealed an increase in pulse frequency at an unusually high rate giving evidence for the first time for a glitch  in an accretion powered pulsar (Galloway 2004). A broad band (0.1-100.0 keV) study of this source was first carried out with {\\em BeppoSAX}. This revealed that the energy spectrum has three components - a Comptonized component, a 0.6 keV blackbody component, and a narrow and weak iron emission line at 6.7 keV (Naik et al. 2006). The absorption column density measured towards this source is low (4.0-5.0$\\times$ 10$^{21}$ atoms cm$^{-2}$).\\\\  Quasi Periodic Oscillations (QPOs) in X-ray binaries are generally thought to be related to the rotation of the inner accretion disk (Paul \\& Rao 1998). Any inhomogeneous matter distribution or blobs of material in the inner disk may result in QPOs in the power spectrum. In the case of accretion powered X-ray pulsars this gives useful information about the interaction between the accretion disk and the neutron star magnetosphere. Black hole X-ray binaries and low magnetic field neutron stars show QPOs over a wide range of frequency from a few Hz to a few hundred Hz. High magnetic field neutron star systems show only low frequency QPOs, in the range 10 mHz upto about 1 Hz. We have investigated the timing properties of the transient X-ray pulsar KS 1947+300 using observations made with the RXTE-PCA and report here the discovery of a transient QPO feature in this source. ",
        "conclusions": "We reported the detection of a low frequency QPO feature in the HMXB pulsar KS 1947+300 at $\\nu_{QPO} = 0.02$ Hz. The power density spectra of accretion powered pulsars consist of narrow peaks corresponding to the spin frequency of the pulsar and its harmonics, accompanied by aperiodic variabilities like broad bumps. QPOs have been detected in 16 accretion powered high magnetic field pulsars. These include many HMXBs and a few LMXBs; both transient  and persistent. The list of transient and persistent sources with QPO features are given in Table 3. There are only three LMXB sources in which QPOs have been detected. Table 3 summarizes the QPO detections in high magnetic field accreting pulsars, giving spin frequency $\\nu_{s}$, the observed QPO frequency $\\nu_{QPO}$, and ratio of the two. The frequency of the QPO feature, detected in KS 1947+300 is lowest among the transient X-ray pulsars.\\\\   Various models have been proposed to explain the sub-Hz QPO features in high mass X-ray binaries. Among the most popular models are the Magnetospheric Beat Frequency Model (MBFM; Alpar and Shaham 1985) and the Keplerian Frequency Model (KFM; van der Klis et al. 1987). In the KFM the QPOs arise from the modulation of the X-rays by inhomogeneities in the accretion disk, at the Keplerian frequency. In the MBFM, the QPO frequency is the difference between the spin frequency and the Keplerian frequency of the inner edge of the accretion disk $\\nu_{QPO}= \\nu_{k}- \\nu_{s}$ (Shibazaki \\& Lamb 1987). Mass flow to the neutron star is expected to be stopped at magnetospheric boundary by centrifugal inhibition of accretion if the Keplerian frequency at the magnetospheric boundary is less than the neutron star spin frequency (Stella et al. 1986). Among the sources listed in Table~\\ref{QPO-list}, the KFM is not applicable for some of the sources including KS 1947+300 as the observed QPO frequency is smaller than the pulse frequency. At larger mass accretion rate, the accretion disk is expected to extend closer to the neutron star, and therefore, in either model, a positive correlation is expected between the QPO centroid frequency and the X-ray intensity (Finger 1998). In some sources, the QPO frequency is found to be quite constant with more than a factor of 10 variation in X-ray luminosity. In some of these sources, it is believed that the inner disk origin may not be applicable while in some other sources there are a variety of other reasons for a lack of correlation (Raichur \\& Paul 2008a,b).\\\\  A positive energy dependence of the QPO rms, similar to that seen in KS 1947+300 is found in some sources (XTE J1858+034, Paul \\& Rao 1998). This favours the MBFM over the KFM, while the QPO rms is independent of the energy in the persistent X-ray pulsar Cen X-3 (Raichur \\& Paul 2008a).  With some exceptions (4U 1626-67; Kaur et al. 2008 but see also Jain et al. 2009), the QPOs are transient phenomena in all types of pulsars. The persistent X-ray pulsars do not show persistent QPOs (Raichur \\& Paul 2008a). In the transient X-ray pulsars that show QPOs, the feature is not detected in all the outbursts. For example the recent outbursts of the source A 0535+262 does not show the flux dependent QPO features that was observed with BATSE during a large outburst in 1994 (Finger et al. 1996). However, if QPOs are present during an outburst, it is usually present throughout the outburst (Finger et al. 1996, Mukherjee et al. 2006). In the case of KS 1947+300, the QPO feature was not seen in 2000 and 2002 data. It appeared very near the end of the 2001 outburst during MJD 52052-52062. We also analysed the 2-20 keV spectrum during the peak of the outburst and the flux measured} is $3\\times$10$^{-9}$ erg cm$^{-2}$ s$^{-1}$. The QPO feature is seen to be present when the X-ray flux had dropped to 1.6\\% of this peak value.  \\\\  Assuming MBFM for KS 1947+300, the Keplerian frequency of the inner disk corresponding to a 0.02 Hz oscillation is $\\nu_{k}= 0.02+0.053 = 0.073$. Thus the radius of the inner accretion disk can be estimated as \\textbf{\\begin{equation} R_{BFM} = \\left({GM}\\over{4 \\pi^{2} (\\nu_{QPO}+\\nu_{s})^{2}}\\right)^{1/3} \\end{equation}} The $R_{BFM}$ obtained for a mass of  $1.4 M_{\\odot}$ is $9.6 \\times 10^{3}$ km. The X-ray flux of KS 1947+300 in the 2-20 keV band is $5.0 \\times 10^{-11}$ erg cm$^{-2}$ $s^{-1}$ during the time of QPO detection. This flux amounts to an X-ray luminosity of $0.6\\times 10^{36}$ erg s$^{-1}$ for a source distance of 10 kpc. Using a correlation between the spin-up torque and the broad band X-ray luminosity of this source a very strong magnetic field of $B = 2.5 \\times 10^{13}$ G was inferred for the neutron star (Tsygankov \\& Lutovinov 2005). Using the above, and canonical values for the stellar radius and mass of 10 km and 1.4 $M_{\\odot}$ respectively, the size of the magnetosphere for a dipole field configuration can be estimated as (Ghosh \\& Lamb 1991)  \\begin{eqnarray}\\nonumber R_{m} = 2.4 \\times 10^3 ({M}/{1.4 M_{\\odot}})^{1/7}\\times ({B}/{2.5 \\times 10^{12} G})^{4/7}\\times \\\\ ({R}/{10^6 cm}) ^{10/7}\\times ({L_{x}}/{3.1 \\times 10^{37} ergs~s^{-1}})^{-2/7} km \\end{eqnarray}   The magnetospheric radius is obtained as $2.76\\times 10^4$ km.   Although the magnetic field value is rather uncertain, such a strong magnetic field implies a magnetospheric radius larger by a factor of 2.8 compared to the inner disk radius at which the QPOs are likely to be produced in the MBFM and by a factor of 2 compared to the corotation radius. This is true even if we assume a 50\\% bolometric correction, 12\\% larger distance (Kiziloglu et al. 2007) and the neutron star mass and radius larger by 20\\% (Lattimer \\& Prakash 2007). If the magnetic field is so strong, the X-ray flux value during the QPO detection indicates that the source should be in propeller regime. The presence of QPOs and the strong pulsations at the low flux value thus suggests that the magnetic field may not be as high as was suggested by Tsygankov \\& Lutovinov (2005). On the other hand, for an X-ray luminosity of $0.6\\times 10^{36}$ erg s$^{-1}$ d$_{10 kpc}^{2}$ if the observed QPOs are produced due to inhomogeneities near the magnetospheric radius, it implies a magnetic field strength of about $B = 4 \\times 10^{12}$ d$_{10 kpc}$ G for the neutron star. The X-ray spectrum of KS 1947+300, however, is devoid of any cyclotron absorption line corresponding to such a magnetic field strength (Naik et al. 2006)."
    },
    "1005/1005.3813_arXiv.txt": {
        "abstract": "{We present \\emph{Herschel}/PACS, MMT/Hectospec and XMM-Newton observations of Abell 1835, one of the brightest X-ray clusters on the sky, and the host of a strong cool core.  Even though Abell 1835 has a prototypically ``relaxed'' X-ray morphology and no signs of ongoing merger activity in strong- and weak-lensing mass maps, it has a complex velocity distribution, suggesting that it is still accreting significant amounts of mass in the form of smaller satellite systems.  Indeed, we find strong dynamical segregation of star-forming dusty galaxies from the optically selected cluster population.  Most \\emph{Herschel} sources are found close to the virial radius of the cluster, and almost a third appear to be embedded within a filament feeding the cluster from the SW. We find that the most luminous infrared galaxies are likely involved in galaxy-galaxy interactions that may have triggered the current phase of star formation. } ",
        "introduction": "Galaxy clusters have long been thought of as particularly active sites of galaxy evolution \\citep{Dressler1980}. Early optical studies suggested that clusters affect stellar populations in infalling substructure mainly by turning off star formation as galaxies encounter the dense cluster environment \\citep{Lewis2002}. However, when these same objects were observed in the infrared by IRAS, ISO and \\emph{Spitzer}, a significant fraction of cluster galaxies was revealed to contain heavily obscured sites of star formation (see e.g. \\citet{Metcalfe2005} and references therein). When this obscured activity is accounted for, clusters emerge as environments where star formation is both quenched and induced, perhaps even at different stages of the same physical process. Some studies have claimed to observe triggered star formation in strongly disrupted, merging clusters \\citep[e.g.][]{Geach2006} but the intrinsic scatter in this correlation is large, with some seemingly relaxed clusters presenting much higher star formation activity than their actively merging counterparts \\citep{Haines2009}.   Many studies in recent years have attempted to separate cluster samples into two distinct populations: merging clusters vs clusters that are dynamically relaxed. The global properties of the latter are expected to be tightly correlated with mass, and therefore directly relatable to statistical predictions from cosmological models.  While there is evidence that some X-ray cluster properties may in fact be bimodal, e.g. with respect to the presence or absence of a cool core \\citep{Sanderson2009}, cosmological simulations have shown that even the most relaxed cool core clusters can still be accreting significant amounts of mass  \\citep{Poole2008}, at a mean rate which is approximately linearly proportional to the mass of the system \\citep{McBride2009}.  Generally, only the most massive of these accretion events (mass ratios $<10:1$) are expected to disturb the dense cluster core. Massive cool core clusters can thus quietly accrete a significant fraction of their total mass without large disturbances to their central ICM,  which makes them arguably ``cleaner\" targets for studies of cluster galaxy evolution.  The Local Cluster Substructure Survey (LoCuSS\\footnote{http://www.sr.bham.ac.uk/locuss}) \\emph{Herschel} key programme \\citep{Smith2010} was designed to probe  dusty star-forming galaxies in the infall regions of a sample of 30 galaxy clusters ($z\\sim0.2$) spanning a wide range of mass and merging histories.   Abell 1835 is the most X-ray luminous cluster in this sample ($L_{\\rm X,bol}=5.32\\pm 0.15 \\times10^{45}$ ergs/s \\citep{Zhang2008}) and the host of a strong cool core \\citep{Peterson2001}. Its strong- and weak-lensing derived mass distributions \\citep{Richard2010,Okabe2010}, and X-ray morphology \\citep{Smith2005} are typical of undisturbed (i.e.\\ non-merging) cool core clusters.  For example, the low substructure mass fraction in its core ($f_{\\rm sub}=0.13\\pm0.01$) suggests that it has grown in mass by $<10\\%$ in the previous few Gyr \\citep{Smith2008}.  In this Letter, we focus on a sample of $100\\mu{\\rm m}$-selected spectroscopically confirmed cluster galaxies in Abell\\,1835. We use the Photodetector Array Camera and Spectrometer (PACS) \\citep{Poglitsch2010} on the \\emph{Herschel} satellite \\citep{Pilbratt2010} to measure far infrared fluxes and from these we estimate total infrared luminosities for these objects. We then use these data to examine the spatial and velocity distribution of obscured star formation sites in Abell 1835. We assume $H_0= 70$km s$^{-1}, \\Omega_M =0.3,\\Omega_\\Lambda=0.7$. All virial masses and radii are stated with respect to the virial overdensity, $\\Delta_{\\rm virial}(z=0.25) =121$ \\citep{Bryan1998}. ",
        "conclusions": "Star formation can be triggered by several different physical processes prevalent in clusters and their infall regions; disentangling these effects is one of the main aims of the LoCuSS \\emph{Herschel} key programme. The most promising candidates, ICM ram pressure, galaxy mergers, and galaxy harassment can all perturb the ISM to produce a new generation of stars, but the relative importance of these processes is a strong function of clustercentric radius (c.f.\\ Fig.~1 in \\citet{Smith2010}). The location of the peak of IR luminosity density just inside the virial radius of Abell 1835 suggests that merging may be driving the bulk of the star formation in these galaxies. This interpretation is corroborated by the large fraction of disturbed morphologies in the infrared luminous sample, and also by our preliminary analysis of their local environment: when compared to a matched sample of galaxies with no detectable infrared emission, infrared bright galaxies are more likely to have close neighbours, although the larger sample of 30 LoCuSS clusters should be used to confirm this correlation.  The $L_{\\rm TIR}$ density profile for Abell 1835 is not spherically symmetric: the brightest galaxies are clustered together in an elongated structure towards the south of the cluster that is dynamically segregated from the rest of the cluster. The structure's K-band morphology and velocity gradient toward the cluster core suggests that these galaxies may be embedded within a filament that is still currently feeding the cluster with new, actively star forming galaxies.  \\emph{XMM} imaging of this field reveals the presence of two groups in the outskirts of Abell 1835 (see contours in Fig.~1), one of which appears to be embedded in the filament that is also feeding the bright LIRGs into the cluster. The group is dominated by a large {(L$_K\\approx 2\\times10^{12}$L$_\\odot, z=0.245$)} elliptical galaxy with no detectable \\emph{Herschel} emission.  The two groups have X-ray luminosities of $4.6\\pm0.6$ and $7.5\\pm0.8$ $\\times10^{42}$ ergs s$^{-1}$ (0.1-2.4 keV), from which we estimate masses, using the weak lensing calibration of the $L_X-M$ relation of \\citet{Leauthaud2010}, of $M_{200} = 4.3$ and $5.8 \\times10^{13} M_\\odot$ respectively. Together these amount to $\\gtrsim$ 1/10th of the mass of the central cluster, reinforcing the view that Abell\\,1835 continues to grow by feeding on smaller systems from the surrounding large scale structure."
    },
    "1005/1005.3022_arXiv.txt": {
        "abstract": "Advances in observational capabilities have ushered in a new era of multi-wavelength, multi-physics probes of galaxy clusters and ambitious surveys are compiling large samples of cluster candidates selected in different ways. We use a high-resolution N-body simulation to study how the influence of large-scale structure in and around clusters causes correlated signals in different physical probes and discuss some implications this has for multi-physics probes of clusters (e.g.~richness, lensing, Compton distortion and velocity dispersion).  We pay particular attention to velocity dispersions, matching galaxies to subhalos which are explicitly tracked in the simulation. We find that not only do halos persist as subhalos when they fall into a larger host, groups of subhalos retain their identity for long periods within larger host halos. The highly anisotropic nature of infall into massive clusters, and their triaxiality, translates into an anisotropic velocity ellipsoid: line-of-sight galaxy velocity dispersions for any individual halo show large variance depending on viewing angle. The orientation of the velocity ellipsoid is correlated with the large-scale structure, and thus velocity outliers correlate with outliers caused by projection in other probes. We quantify this orientation uncertainty and give illustrative examples. Such a large variance suggests that velocity dispersion estimators will work better in an ensemble sense than for any individual cluster, which may inform strategies for obtaining redshifts of cluster members. We similarly find that the ability of substructure indicators to find kinematic substructures is highly viewing angle dependent. While groups of subhalos which merge with a larger host halo can retain their identity for many Gyr, they are only sporadically picked up by substructure indicators.  We discuss the effects of correlated scatter on scaling relations estimated through stacking, both analytically and in the simulations, showing that the strong correlation of measures with mass and the large scatter in mass at fixed observable mitigate line-of-sight projections. ",
        "introduction": "Galaxy clusters form the high-mass tail of hierarchical structure formation and are of interest for constraining cosmological parameters, understanding large scale structure, as extreme environments for galaxy formation and as objects hosting unique astrophysical phenomena. While first discovered as concentrations of galaxies \\citep{Abe58,Zwi66}, they are now also routinely found as luminous, extended X-ray sources \\citep{Sch78,McH78,Boh00,Boh01}, as peaks in the shear field \\citep{Wit06} and as ``holes'' in the microwave sky \\citep{SPT}. To mitigate the systematic errors associated with each individual method and to provide a more complete understanding of clusters, multi-wavelength studies have become increasingly common.  Each waveband adds knowledge about clusters. However, we might expect there to be significant correlations between effects in different methods both because the intrinsic properties they measure depend on e.g.~cluster size but also because they are similarly affected by the complex environment surrounding clusters.  \\begin{figure} \\begin{center} \\resizebox{4in}{!}{\\includegraphics{greyplot.eps}} \\end{center} \\caption{The velocity field traces the filamentary large-scale structure. The grey scale shows the logarithm of the projected dark matter density in a $5\\,h^{-1}$Mpc thick slice around a cluster of mass $5\\times 10^{14}\\,h^{-1}M_\\odot$ at $z=0$. The slice is oriented to contain the halo particle velocity eigenvectors with the largest ($x$) and smallest ($y$) eigenvalues. The dashed circle shows the region within $r_{200c}$. Note that within the cluster the velocities trace the elongation of the matter, as required.  In addition the velocities at larger distances trace the filamentary structure in the larger scale environment.  Thus we expect that velocity anisotropy will be correlated with density anisotropy.} \\label{fig:dotplot} \\end{figure}  Roughly speaking, both the hot gas and galaxies in clusters trace the dark matter which dominates the potential. We can approximate the clusters as self-similar and isothermal, with a temperature $T\\propto M^{2/3}$, a velocity dispersion $\\sigma^2\\propto T$, and richness $N\\propto M$ for sufficiently massive halos \\citep{Kai86}. Consider a small region near the cluster: lensing measures the sum of all of the mass, richness all of the mass in halos above some threshold and Compton (or SZ) distortion \\citep[][SZ]{SunZel72} all of the mass in halos weighted by $M^{2/3}$. The signal in each probe depends on the mass within the virialized region of the cluster, the mass near the cluster but outside the virial region and (uncorrelated) mass at larger distance along the line-of-sight. Lensing and Compton distortion measures provide little line-of-sight resolution. The degree of projection involved in a richness estimator depends on how well the galaxy distances are known (e.g.~using photometric or spectroscopic redshifts).  Line-of-sight galaxy velocities in principle provide a measure of the potential well depth or mass and offer the possibility of breaking line-of-sight projection. However, the velocity field traces the density field, and can be correlated with line-of-sight projection due to the filamentary nature of mass accretion onto massive halos (see Fig.~\\ref{fig:dotplot} which gives an example of this effect in our simulations). The velocities of cluster galaxies can retain this large scale anisotropy \\citep[see also][for studies of dark matter velocities]{Tor97,KasEvr05}. Thus it is easy to imagine that line-of-sight velocity dispersion could be correlated with filamentary material which can bias individual cluster measurements in e.g.~richness, lensing or Compton distortion.  We would like to investigate how the complex structure of the cosmic web of material near clusters leads to correlations in individual cluster observables, and the implications that this has for these four probes of clusters. This shared dependence, not only on cluster properties but also on cluster environment, can introduce additional subtleties when methods are combined. For example, an often used approach is to ``stack'' clusters on the basis of one observed property $X$ (e.g.~richness), and then look for correlations between two other properties, $Y$ and $Z$. Clearly, it is very important to understand the joint distribution $P(X,Y,Z)$ and the degree of correlation between scatter in $X$, $Y$ or $Z$.  In this paper we use N-body numerical simulations with subhalos (which we identify with galaxies) to study properties of cluster galaxy kinematics and the relation of the scatters in velocity dispersion, Compton distortion, lensing, and optical richness\\footnote{We do not address X-ray emission in this paper.} generated by nearby large-scale structure. Details of our numerical simulations and methods for finding subhalos are given in \\S\\ref{sec:sims}.  We describe how the mock richness, lensing and Compton distortion observations are constructed in \\S\\ref{sec:mock}. Readers interested in the results may skip to \\S\\ref{sec:intrinsic} where we discuss the intrinsic properties of our massive halos and their subhalo (galaxy) populations and \\S\\ref{sec:interlopers} where we discuss measurements of galaxy kinematics in the presence of interlopers and \\S\\ref{sec:observed} where we discuss the correlations between different observables.  The effects of the cosmic web, and in particular projection effects, have been long-time concerns for optical cluster finding \\citep[e.g.][]{Abe58,Dal92,Lum92,vHaarlem97,Whi99}, measuring Compton distortion \\citep[e.g.][]{WhiHerSpr02,HolMcCBab07,Hal07,ShaHolBod08}, or interpreting weak lensing maps \\citep[e.g.][]{RebBar99,MetWhiLok01,Hoe01,PutWhi05,Men10}. Correlations between scatters induced by common projection effects were noted in \\citet{CohWhi09}. \\citet{Cen97} did an early simulation study of projection on several of the indicators we consider here including richness, velocity dispersions and lensing, and measured substructure using dark matter particles. For cluster kinematics in particular, the velocity dispersion properties of dark matter particles and their relation to the cosmic web were studied in \\citet{Tor97,KasEvr05} and \\citet{Biv06} noted that filamentary inflow was expected to affect measured velocity dispersions. Our simulations have enough dynamic range that we can simulate a representative cosmological volume, including the neighboring large-scale structure and cosmic web, while simultaneously resolving and tracking the subhalos which we believe are galactic hosts.  This preserves any correlations between subhalo properties and halo orientation or cosmic web, and coherence between subhalo populations which fell in as part of a group. We emphasize the effect that anisotropy in galaxy kinematics has on line-of-sight velocity dispersion or virial mass estimators of cluster mass and discuss how this scatter compares to (and correlates with) other measures of cluster mass which are sensitive to the cluster's environment. This forms a partial extension of the work of \\citet{Sta10}, who discussed the correlations in scatters of a large number of different intrinsic (rather than projected) cluster quantities, including X-ray. ",
        "conclusions": "Advances in observational capabilities and a new generation of wide-field surveys have led to an explosion in multi-wavelength samples of galaxy clusters.  By studying a cluster in many different wavebands, and from many different approaches, we can obtain complementary information about the physical state of the clusters and mitigate the systematic errors in any single measurement.  Combining different measurements of cluster properties has to be done carefully however, because the environment in which clusters form leads to features which can be correlated across methods.  As the correlation is not perfect, such a combination will provide improvements over any individual method if done correctly.  In this paper we have used high resolution N-body simulations of a cosmological volume to study how the large-scale environment of clusters leads to correlated scatter in measures of cluster size, specifically those based upon richness, Compton distortion, lensing or velocity dispersions. Our simulation has enough force and mass resolution to track the subhalos which are expected to host galaxies, allowing us to study dynamical probes of the cluster with realistic samples incorporating a hierarchy of substructures and retaining correlations between subhalo positions, velocities and environment. For this reason we paid particular attention to dynamical probes of cluster size.  As might be expected in hierarchical structure formation scenarios, groups of subhalos retain their identity for long periods within larger host halos. This leads to a ``lack of virialization'' which implies that substructures can thus behave quite coherently in phase space. The highly anisotropic nature of infall into massive clusters, and their triaxiality, means that line-of-sight galaxy velocity dispersions (or virial masses) for any individual halo show large variance depending on viewing angle. This suggests that dispersion-based mass estimators will work better in an ensemble sense than for any individual cluster and that obtaining more than tens of redshifts for any given cluster will not reduce the inferred mass error.  We discussed the effect interloper galaxies, and their removal, has on kinematic measurements and compared different schemes for interloper removal. Results were presented both for individual clusters and for a ``stacked'' ensemble cluster.  All of our simulated clusters contain highly evident substructure, with groups of subhalos which fall in together moving in a coherent fashion for several Gyr.  However standard substructure indicators frequently miss this substructure, and often give very different answers for a single object viewed from different directions. These results suggest caution when interpreting lack of observed substructure in the galaxy distribution as evidence for a dynamically relaxed, steady-state object (e.g.~justifying the use of the virial theorem or Jeans analysis without the time derivative). A cluster can be undergoing substantial mass accretion, i.e.~be far from steady state, and still not be seen to have substructure along some lines of sight. The viewing angle dependence also complicates inferences about incidence of dynamical evolution of cluster galaxies from observed interactions of subclusters within the cluster identified through substructure finding techniques.  Finally we note that many observational probes of clusters suffer from projection effects, and that these are exacerbated by the dense, active and anisotropic environments surrounding these massive objects. We found increased nearby massive and less massive halos, and substructure, when two of our measures (richness, lensing, Compton distortion and velocity dispersions) simultaneously had large outliers in predicted mass.  The converse was not always true, scatter in environment or the measurement of substructure did not necessarily imply large outliers.  Since the orientation of the velocity ellipsoid is correlated with the large-scale structure, velocity outliers also correlate with projection induced outliers.  For many cases the same structure causes scatter in different observations: such scatters can be substantially correlated and this correlation needs to be properly incorporated when combining measurements.  \\bigskip  We would like to thank J. Bullock, A. Evrard, B. Gerke, H.Hoekstra, E. Rozo, E. Rykoff, C. Stubbs, A. Wetzel and  A. Zabludoff for conversations and A. Evrard and A. Wetzel for comments on the draft. We thank A.~Zabludoff for suggesting we consider the ``virial mass'' $R_h\\sigma^2$ in addition to $\\sigma^3$. M.W.~thanks Charlie Conroy and James Gunn for useful conversations and collaborations about the stellar population synthesis technique. We thank the referee, Andrea Biviano, for a constructive report and helpful suggestions. The simulations used in this paper were performed at the National Energy Research Scientific Computing Center and the Laboratory Research Computing project at Lawrence Berkeley National Laboratory. M.W. was supported by the DoE and NASA. R.S. was supported by the University of California Education Abroad Program. J.D.C. thanks LBNL for travel support to the SNOWCLUSTER meeting and thanks its organizers and participants for the opportunity to present this work and for their suggestions and questions.  \\appendix"
    },
    "1005/1005.1211_arXiv.txt": {
        "abstract": "{The dust destruction timescales in the cores of clusters of galaxies are relatively short given their high central gas densities. However, substantial mid-infrared and sub-mm emission has been detected in many brightest cluster galaxies. In this letter we present {\\it Herschel} \\rm PACS and SPIRE photometry of the brightest cluster galaxy in three strong cooling flow clusters, A1068, A2597 and Zw3146. This photometry indicates that a substantial mass of cold dust is present ($>3\\times 10^7$~M$_\\odot$) at temperatures significantly lower (20--28~K) than previously thought based on limited MIR and/or sub-mm results. The mass and temperature of the dust appear to match those of the cold gas traced by CO with a gas-to-dust ratio of 80--120.}   \\date{Received 30 March 2010/Accepted} ",
        "introduction": "The cores of cluster of galaxies are very energetic regions with a high X-ray emissivity, particle density, cosmic ray flux, stellar density and AGN radiation. In this very hostile environment any dust grains are unlikely to survive for more than a few million years due to the action of collisional sputtering (Dwek \\& Arendt 1992) unless they are shielded (Fabian et al. 1994). It is therefore somewhat surprising to find that dust continuum emission from the brightest cluster galaxies in the most rapidly cooling clusters being detected at sub-mm and MIR wavelengths (Edge et al. 1999, Egami et al. 2006, O'Dea et al. 2008). The presence of cold molecular gas (Edge 2001, Salom\\'e \\& Combes 2003) and dust absorption in HST imaging (McNamara et al 1996) implies that the dust continuum traces a substantial, cold component to the ISM in these massive elliptical galaxies. However, the origin of the dust and how it is shielded are still poorly understood.  The limitations with the current observations of dust emission make it difficult to establish an unambiguous dust mass as they do not sample over the peak of the dust emission in the FIR. The unprecedented sensitivity of {\\it Herschel} \\rm (Pilbratt et al. 2010) to FIR continuum offers the opportunity to accurately constrain the full FIR spectrum of the dust emission in cluster cores. The authors were awarded 140~hours of time in an Open Time Key Project (PI Edge) to investigate the FIR line and continuum properties of a sample of 11 brightest cluster galaxies (BCGs) in well-studied cooling flow clusters selected on the basis of optical emission line and X-ray properties. The full goals of the project are to observe at least five atomic cooling lines for each object that cover a range in density and temperature behaviour and obtain a fully sampled FIR spectral energy distribution. In this paper we present the Photodetector Array Camera \\& Spectrometer (PACS, Poglitsch et al. 2010) and Spectral and Photometric Imaging REceiver (SPIRE, Griffin et al. 2010) photometry for the three targets observed in the Science Demonstration Phase (SDP), Abell 1068 ($z=0.1386$), Abell 2597 ($z=0.0821$) and Zw3146 ($z=0.2906$). In a parallel paper (Edge et al. 2010), we present the FIR spectroscopy for the first two of these clusters.  The three clusters observed have quite contrasting multiwavelength properties. Abell 1068 and Zw3146 both have strong MIR emission (O'Dea et al. 2008, Egami et al. 2006) with a relatively bright CO detection (Edge 2001) and a weak central radio source (McNamara et al. 2004). A1068 lies just below the luminosity threshold of a ULIRG (10$^{12}$~L$_\\odot$) and exhibits some contribution from an AGN (Crawford et al. 1999, O'Dea et al. 2008). On the other hand, Abell 2597 is a relatively weak MIR source (Donahue et al. 2007) with a weak CO detection (Salome, priv. comm.) and a powerful central radio source (Sarazin et al. 1995). The implied FIR luminosity of A2597 is a factor of around 30 below that of A1068 and, in addition, the fractional contribution from an AGN in the MIR is also lower in A2597. ",
        "conclusions": "Our initial {\\it Herschel} results confirm the presence of the striking dust emission peak expected from the observations at sub-mm (Edge et al. 1999, Chapman et al. 2002) and MIR (Egami et al. 2006, O'Dea et al. 2008).  The star formation rates derived from the full-sampled FIR SED are comparable to those derived from {\\it Spitzer} 24$\\mu$m fluxes apart from  A1068, which has the strongest contribution from an AGN so hot dust dominates to the 24$\\mu$m flux. However, in the sub-mm the contribution from the radio continuum from an active nucleus must be correctly accounted for before any dust mass can be estimated from the 850$\\mu$m flux. In the case of A2597 here and A2390 in Edge et al. (1999), the presence of a powerful radio source appears to contribute to the SCUBA 850$\\mu$m flux.  While it is difficult to draw any general conclusions from just three BCGs, we note with interest that the ratio of dust mass to CO-derived gas mass is constistent for all three within a factor of five. If the dust were mostly generated through dust ejection from evolved stars then the dust mass should closely correlate with the total stellar mass. However, our three galaxies have very similar optical/NIR absolute magnitudes. So, unless the ejected dust were ``captured'' by the cold gas clouds  protecting it from X-ray sputtering, this suggests that the apparent correlation between the molecular gas and dust masses arises from a direct connection between the gas reservoir and star formation.   These results are a limited example of those to come in the very near future from {\\it Herschel} \\rm as there are two other Open Time Key Projects (PI Egami and Smith) that are targetting a total of 70 clusters that cover a broad range of BCG properties so the wider context of these initial observations can be determined. In particular, the amount of dust present in more quiescent BCGs and other massive cluster ellipticals will be important in assessing how much of the dust seen in cool core BCGs orginates from the underlying stellar population.    \\begin{figure} \\centerline{\\includegraphics[width=5.7cm,angle=90]{14572f1a.ps}} \\centerline{\\includegraphics[width=5.7cm,angle=90]{14572f1b.ps}} \\centerline{\\includegraphics[width=5.7cm,angle=90]{14572f1c.ps}} \\caption{ Spectral energy distributions for A1068 (top), A2597 (middle) and Zw3146 (bottom) including {\\it Herschel} \\ PACS/SPIRE (blue/red symbols), {\\it Spitzer}, Radio, NIR photometry from 2MASS and optical photometry from SDSS. % To account for absolute flux uncertainties we have set the following errors on the fluxes derived from the various instruments (unless the quoted error is larger than this); PACS BS/BL 10\\%, PACS R 20\\%, SPIRE 15\\%, {\\it Spitzer} 10\\%, SCUBA 450 $\\mu$m 30\\% and SCUBA 850 $\\mu$m 20\\%. The model fit to the sub-mm/FIR/MIR data is shown by the black solid line. Only filled symbols have been used in the fit. The two modified blackbodies making up the model are shown by the black long dash and dash-dot lines. For A2597 we also show two VLBI measurements (black crosses) of the BCG core at 1.3 and 5 GHz (Taylor et al. 1999). These points show that the BCG has a strong, inverted radio core.} \\end{figure}"
    },
    "1005/1005.2413_arXiv.txt": {
        "abstract": "We measure the luminosity and color dependence of galaxy clustering in the largest-ever galaxy redshift survey, the main galaxy sample of the Sloan Digital Sky Survey (SDSS) Seventh Data Release (DR7). We focus on the projected correlation function $\\wrp$ of volume-limited samples, extracted from the parent sample of $\\sim 700,000$ galaxies over $8000$ deg$^2$, extending up to redshift of $0.25$. We interpret our measurements using halo occupation distribution (HOD) modeling assuming a $\\Lambda$CDM cosmology (inflationary cold dark matter with a cosmological constant). The amplitude of $\\wrp$ grows slowly with luminosity for $L<L_*$ and increases sharply at higher luminosities, with a large-scale bias factor $b(>L) \\times (\\sigma_8/0.8) = 1.06+0.21(L/L_*)^{1.12}$, where $L$ is the sample luminosity threshold. At fixed luminosity, redder galaxies exhibit a higher amplitude and steeper correlation function, a steady trend that runs through the ``blue cloud'' and ``green valley'' and continues across the ``red sequence.''  The cross-correlation of red and blue galaxies is close to the geometric mean of their auto-correlations, dropping slightly below at $r_p < 1\\hmpc$. The luminosity trends for the red and blue galaxy populations separately are strikingly different. Blue galaxies show a slow but steady increase of clustering strength with luminosity, with nearly constant shape of $\\wrp$.  The large-scale clustering of red galaxies shows little luminosity dependence until a sharp increase at $L>4L_*$, but the lowest luminosity red galaxies ($0.04-0.25L_*$) show very strong clustering on small scales ($r_p<2\\hmpc$). Most of the observed trends can be naturally understood within the $\\Lambda$CDM+HOD framework. The growth of $\\wrp$ for higher luminosity galaxies reflects an overall shift in the mass scale of their host dark matter halos, in particular an increase in the minimum host halo mass $M_{\\rm min}$. The mass at which a halo has, on average, one satellite galaxy brighter than $L$ is $M_1 \\approx 17\\Mmin(L)$ over most of the luminosity range, with a smaller ratio above $L_*$.  The growth and steepening of $\\wrp$ for redder galaxies reflects the increasing fraction of galaxies that are satellite systems in high mass halos instead of central systems in low mass halos, a trend that is especially marked at low luminosities. Our extensive measurements, provided in tabular form, will allow detailed tests of theoretical models of galaxy formation, a firm grounding of semi-empirical models of the galaxy population, and new constraints on cosmological parameters from combining real-space galaxy clustering with mass-sensitive statistics such as redshift-space distortions, cluster mass-to-light ratios, and galaxy-galaxy lensing. ",
        "introduction": "\\label{sec:intro}  Three-dimensional maps of the large-scale distribution of galaxies reveal a rich network of filaments and sheets, punctuated by dense clusters and interleaved with low density tunnels and bubbles \\citep{gregory78,kirshner81,davis82,giovanelli86,geller89,shectman96,colless01}. Different classes of galaxies trace this structure differently, with early type galaxies residing preferentially in rich groups and clusters and late type galaxies residing preferentially in the filaments and walls; this segregation of clustering was already evident in two-dimensional studies as early as \\cite{hubble36}. Galaxy surveys map the distribution of visible baryons, but a combination of observational and theoretical arguments, beginning with Zwicky (\\citeyear{zwicky33}, \\citeyear{zwicky37}), show that the galaxies trace an underlying network of invisible, gravitationally dominant dark matter. In this paper, we measure the clustering of galaxies as a function of luminosity and color in the largest galaxy redshift survey to date, the main galaxy sample \\citep{strauss02} of the Seventh Data Release (DR7; \\citealt{abazajian09}) of the Sloan Digital Sky Survey (SDSS; \\citealt{york00}). Our primary tool is the two-point correlation function $\\xi(r)$, which provides a simple, robust, and informative measure of galaxy clustering (e.g., \\citealt{peebles80}). More specifically, we focus on the projected correlation function $\\wrp$, which  integrates out redshift-space distortions caused by galaxy peculiar velocities \\citep{davis83}. By modeling our measurements in the context of the $\\Lambda$CDM cosmological framework (inflationary cold dark matter with a cosmological constant), we infer the relation between different classes of galaxies and the underlying distribution of dark matter, providing fundamental tests for theories of galaxy formation.   Over the last few decades, a variety of ``local'' clustering studies have established an increasingly refined and quantitative characterization of the dependence of galaxy clustering on luminosity, morphology, color, and spectral type (e.g., \\citealt{davis76,davis88,hamilton88,alimi88,valls89,loveday95,benoist96,guzzo97,willmer98,brown00,norberg01,norberg02,zehavi02,budavari03,madgwick03,zehavi05b,li06,swanson08,loh09}). Luminous galaxies generally cluster more strongly than faint galaxies, reflecting their tendency to reside in denser environments. Galaxies with bulge-dominated morphologies, red colors, or spectral types indicating old stellar populations also exhibit stronger clustering and a preference for dense environments. Significant progress has also been made in recent years in measuring galaxy clustering at intermediate and high redshifts (e.g., \\citealt{brown03,daddi03,adelberger05b,ouchi05,lee05,phleps05,coil06,coil07,meneux07,meneux09,abbas10}).  Cosmological inferences from galaxy clustering measurements are complicated by the existence of galaxy bias, the difference between the distribution of galaxies and that of the underlying dark matter. While the gravitational clustering of dark matter from specified initial conditions can be computed reliably with cosmological $N$-body simulations, the detailed physics of galaxy formation --- gas cooling, star formation, and the feedback effects of star formation and black hole accretion --- is only partly understood, so galaxy bias cannot be predicted robustly from first principles.  Cosmological parameter studies must adopt a mathematical description of galaxy bias and marginalize over its uncertain parameters.  This procedure is most straightforward at large scales, where the effects of bias are expected to be simple, i.e., a scale-independent amplification of the matter $\\xi(r)$ \\citep{kaiser84,BBKS86,coles93,fry93,mann98,scherrer98,narayanan00}. Conversely, for a specified cosmological model, one can constrain detailed descriptions of galaxy bias and thus gain insights into galaxy formation physics.  In the cold dark matter scenario \\citep{peebles82,blumenthal84}, which is now supported by a wide range of observational evidence (e.g., \\citealt{dunkley09,reid10}), galaxies form and reside in extended dark matter halos. The existence of such halos is well established by studies of spiral galaxy rotation curves (e.g., \\citealt{rubin78,persic96,verheijen01}) and the stellar dynamics (e.g., \\citealt{gerhard01}) and gravitational lensing (e.g., \\citealt{bolton08}) of elliptical galaxies. Studies of weak lensing and satellite galaxies show that the halos of luminous galaxies extend to hundreds of kpc, where they join smoothly onto the larger scale distribution of dark matter (e.g., \\citealt{zaritsky94,fischer00,prada03,mandelbaum06}). The formation of dark matter halos is dominated by gravity and can be well predicted for a given cosmology from high-resolution numerical simulations and analytic models. Dark matter halos thus become the natural bridge for connecting the galaxy distribution and the matter distribution.  In recent years, the theoretical understanding of galaxy clustering has been enhanced through development of the Halo Occupation Distribution (HOD) framework (e.g., \\citealt{jing98a,ma00,peacock00,seljak00,scoccimarro01,berlind02,cooray02, yang03,kravtsov04,zheng05}). The HOD formalism describes the ``bias'' relation between galaxies and mass at the level of individual dark matter halos, in terms of the probability distribution that a halo of virial mass $M_h$ contains $N$ galaxies of a given type, together with prescriptions for the relative spatial and velocity bias of galaxies and dark matter within virialized halos. The combination of a cosmological model and a fully specified HOD can predict any galaxy clustering statistic on any scale, allowing integrated constraints from many observations. For an assumed cosmological model and a parameterized form of the HOD motivated by contemporary theories of galaxy formation (e.g., \\citealt{kauffmann97,kauffmann99,benson00,berlind03,kravtsov04,zheng05,conroy06}), measurements of $\\wrp$ are already highly constraining, and HOD modeling transforms data on galaxy pair counts into a physical relation between galaxies and dark matter halos. HOD modeling has been applied to interpret clustering data from a number of surveys at low and high redshifts (e.g., \\citealt{jing98b,jing02,bullock02,moustakas02,bosch03a, magliocchetti03,yan03,zheng04,yang05b,zehavi05b,cooray05b,hamana05, lee05,lee08}; Zheng, Coil \\& Zehavi 2007; \\citealt{white07,blake08,brown08,quadri08,wake08,kim09,zheng09,ross10}).  In principle, a complete model of galaxy bias might need to allow for the possibility that the average galaxy content of halos depends on large-scale environment as well as halo mass, since halo concentrations and assembly histories show some environmental correlations \\citep{sheth04,gao05,harker06,wechsler06,zhu06,croton07,jing07,wetzel07,dalal08,zu08}. However, studies assuming an environment-independent HOD have proven successful at explaining galaxy clustering in different density regimes (\\citealt{abbas06,abbas07,tinker08,skibba09}; see also \\citealt{blanton06,blanton07}), and theoretical models predict only a small impact of such ``halo assembly bias'' on galaxy clustering statistics for mass- or luminosity-thresholded samples \\citep{yoo06,croton07,zu08}.  The present paper builds upon our investigation of galaxy correlations in early SDSS redshift data \\citep{zehavi02}, our use of HOD modeling to interpret deviations from a power-law in the galaxy two-point correlation function \\citep{zehavi04}, and, especially, our earlier investigation of luminosity and color dependence of the galaxy correlation function in a sample of about 200,000 SDSS galaxies (Zehavi et al.\\ 2005b, hereafter Z05). Here we take advantage of the final SDSS galaxy sample --- roughly three times more galaxies once appropriate cuts are applied --- and advances in HOD modeling methods to obtain higher precision measurements and tighter, more informative constraints on galaxy-halo relations. This study complements correlation function measurements and HOD models of the SDSS Luminous Red Galaxy (LRG) sample \\citep{eisenstein01,eisenstein05b,zehavi05a,zheng09,kazin10,watson10}, which probes the most luminous galaxies out to redshift $z\\approx 0.45$ \\citep{eisenstein01}. We focus our analysis on volume-limited samples of well defined galaxy classes, which allows us to construct HOD models with a small number of free parameters to interpret the measurements for each class. This approach complements other analyses of the SDSS main galaxy sample that measure luminosity or stellar mass-weighted correlation functions \\citep{li09,li10} or use marked correlation functions to quantify luminosity and color dependence \\citep{skibba06,skibbasheth09}. These analyses typically yield smaller error bars because they use more sample galaxies for the measurement, but they require a more complete global description of the galaxy population to model the results. There are many parallels between our program and the one pursued by van den Bosch, Mo, Yang, and their collaborators (e.g., papers cited above and \\citealt{bosch03b,weinmann06,bosch07,yang08,more09}), though they have largely focused on analysis of group catalogs \\citep{yang05a,yang07} rather than detailed fitting of the correlation function.  The two approaches yield qualitatively similar results (e.g., Z05; \\citealt{yang05b,zheng07,yang08}).  Our correlation function measurements provide basic empirical characterizations of large-scale structure at low redshift ($z<0.25$), and the luminosity and color dependence of these correlation functions can test predictions from hydrodynamic cosmological simulations (e.g., \\citealt{pearce01,weinberg04}) or semi-analytic models (e.g., \\citealt{kang05,croton06,bower06}). The derived HOD constraints provide informative tests of galaxy formation models, a low redshift baseline for evolutionary studies (e.g., \\citealt{zheng07,brown08}), and a description that can be used to create realistic mock catalogs from simulations (\\citealt{scoccimarro02,wechsler04,eisenstein05b,skibbasheth09}; McBride et al., in preparation). As discussed extensively by \\cite{zheng07a}, the HOD formalism can also be used in cosmological parameter determinations, allowing marginalization over the parameters of a bias prescription that applies to a wide range of clustering statistics from the linear to the highly non-linear regime.  Combinations of spatial clustering statistics and dynamically sensitive measures (such as galaxy-galaxy lensing, redshift-space distortions, or group and cluster mass-to-light ratios) can break the main degeneracies between cosmological parameters and galaxy bias.  A number of papers have implemented variants of this approach to constrain the matter density parameter $\\Omega_m$ and the amplitude of matter clustering $\\sigma_8$ \\citep{bosch03b,abazajian05,tinker05,bosch07,cacciato09,rozo10}. These analyses argued for a significant downward revision of the WMAP1 values of $\\Omega_m$ and/or $\\sigma_8$, anticipating the parameter changes that occurred with WMAP3 (see also \\citealt{vale06}, who reached a similar conclusion by a related method). Our correlation function measurements are providing essential constraints for such analyses using the SDSS DR7 data set.  The paper is organized as follows. Section~\\ref{sec:obs} describes the SDSS data and the methods we use to measure galaxy clustering and to interpret it via HOD modeling.  In \\S\\ref{sec:lum} we present results on the luminosity dependence of $\\wrp$ and its implications for HOD models. In \\S~\\ref{sec:color} we examine the dependence of clustering on galaxy color, including cross-correlations between red and blue galaxy samples, and we investigate the luminosity dependence for red and blue galaxies separately. Section~\\ref{sec:conclusion} summarizes our results. Appendix A discusses some technical issues relating to predictions of the galaxy cross-correlation function. Appendix B illustrates the robustness of our measurements to different systematics. Appendix C presents in tabular form the $\\wrp$ measurements for most of the samples discussed in the paper. ",
        "conclusions": "\\label{sec:conclusion}  The SDSS galaxy redshift survey allows high-precision clustering measurements for a broadly selected galaxy sample with extensive, high quality photometric information.  We have examined the luminosity and color dependence of the galaxy correlation function in the DR7 main galaxy sample, which includes approximately $700,000$ galaxies over $8000$ deg$^2$, with a median redshift of $\\sim 0.1$. This is the largest sample used to date for such studies, by a factor of several.  Furthermore, the DR7 main galaxy sample is likely to remain the definitive low redshift galaxy survey for many years; other ongoing and planned surveys, including the BOSS survey of LRGs in SDSS-III \\citep{schlegel09}, will probe larger volumes and higher redshifts, but they will not target a wide range of galaxy types in the present-day universe.  Our analysis focuses on the projected auto-correlation functions calculated for volume-limited samples defined by luminosity and color cuts, with measurements tabulated in Appendix~\\ref{sec:measurements}. We use HOD modeling to interpret these measurements in terms of the relation between galaxies and dark matter halos, assuming a $\\Lambda$CDM cosmological model with $\\Omega_m=1-\\Omega_\\Lambda=0.25$, $\\Omega_b=0.045$, $h=0.7$, $n_s=0.95$, and $\\sigma_8=0.8$.  The amplitude of $\\wrp$ increases with increasing galaxy luminosity, slowly for $L<L_*$ and rapidly for $L>L_*$, where $L_*$ corresponds to $M_r=-20.44$ (\\citealt{blanton03c}; we quote absolute magnitudes for $h=1$ throughout the paper).  For $L \\leq L_*$, $\\wrp$ is reasonably % described by a power-law at $r_p<10\\hmpc$, while brighter samples show clear and increasingly strong inflections at $r_p \\approx 1-3\\hmpc$. We find similar trends for samples defined by luminosity bins and by luminosity thresholds.  The large-scale bias factor of luminosity-threshold samples is well described by the fitting formula $b_g(>L) = 1.06 + 0.21(L/L_*)^{1.12}$.  For luminosity-bin samples, we find $b_g(L) = 0.97+ 0.17(L/L_*)^{1.04}$, similar to the luminosity dependence found by \\cite{norberg01} for $b_J$-selected galaxies in the 2dFGRS.  At fixed luminosity, the redshift-space correlation function of red galaxies exhibits stronger ``finger-of-God'' distortions than that of blue galaxies, while the blue galaxies exhibit stronger large-scale, coherent flow distortions. The projected correlation function of red galaxies is steeper and higher in amplitude.  The cross-correlation of red and blue galaxies is equal to the geometric mean of the auto-correlation functions on large scales, but it falls slightly below the geometric mean for $r_p \\la 1\\hmpc$.  Adopting fine color bins, we find a continuous trend of clustering with color: the bluest galaxies have a shallow, low-amplitude correlation function, the clustering of ``green valley'' galaxies is intermediate between that of blue and red galaxies, and the reddest galaxies have a (slightly) steeper correlation function than galaxies that trace the ridge of the red sequence. We present detailed results for the $-20<M_r<-19$ luminosity bin, but we find similar trends in other bins where our statistics are good enough to measure them.  The luminosity dependence of clustering for the red and blue populations is strikingly different.  For blue galaxies, the amplitude of $\\wrp$ increases slowly but steadily with luminosity over the range $M_r=-18$ to $M_r=-22$, with nearly constant shape. For red galaxies, there are only weak luminosity trends over the range $-22<M_r<-19$.  The $-23<M_r<-22$ galaxies have a much higher correlation amplitude and a strong break in $\\wrp$ at $r_p\\approx 2\\hmpc$.  Most remarkably, the small-scale ($r_p < 2\\hmpc$) correlation function of the $-19<M_r<-18$ red galaxies is equal to that of the $-23<M_r<-22$ red galaxies, a factor of $2-3$ higher than that of intermediate luminosity red galaxies.  Red galaxies with $-18<M_r<-17$ show even stronger small-scale clustering, though our survey volume for such low luminosity systems is small.  Our HOD modeling shows that these varied trends in the amplitude and shape of $\\wrp$ can, for the most part, be well explained by the combination of $\\Lambda$CDM cosmology and physically plausible recipes for the relation between galaxies and dark matter halos.  The luminosity dependence of $\\wrp$ arises from an overall shift in the mass scale of the mean occupation function $\\langle N(M_h)\\rangle$.  The halo mass $\\Mmin$ for hosting central galaxies of luminosity $L$ rises with luminosity. Correspondingly, the central galaxy luminosity increases with halo mass as $L/L_*=A(\\Mmin/M_t)^{\\alpha_M}\\exp(-M_t/\\Mmin+1)$, where $A=0.32$, $M_t=3.08\\times 10^{11}\\hMsun$, $\\alpha_M=0.264$, and $M_{\\rm min}$ is the halo mass at which the median luminosity of central galaxies is $L$. The mass $M_1$ at which halos host an average of one satellite above luminosity $L$ follows a similar trend: we find $M_1\\approx 17 M_{\\rm min}$ over most of our luminosity range, with a smaller factor at the highest luminosities. We find substantial scatter ($\\approx 0.3$ dex) between halo mass and central galaxy luminosity for $L>L_*$, while fits for lower luminosities are consistent with little or no scatter.  The color dependence of $\\wrp$ at fixed luminosity can be explained well by a change in the relative fractions of central and satellite galaxies.  In our best-fit models of the $-20<M_r<-19$ bin, for example, the satellite fraction rises steadily from 15\\% for the bluest galaxies to 75\\% for the reddest galaxies. Increasing the satellite fraction increases the large-scale bias factor by placing more galaxies in high mass halos, and it produces a steeper correlation function with a stronger inflection by boosting the one-halo term relative to the two-halo term.  A modest offset in the halo mass scale for {\\it central} red and blue galaxies is physically plausible, but our models are able to fit the main observed trends without such offsets, and it is unlikely that the central-galaxy mass scale can be the primary driver of the observed color trends in $\\wrp$.  Differences in satellite fractions largely explain the different luminosity dependence of $\\wrp$ for red and blue galaxies. However, within our standard parameterization we are unable to find a statistically acceptable fit to the clustering of the red $-19 < M_r < -18$ galaxies.  After adjusting the model to allow red satellites only in relatively high mass halos ($M_h > 10^{13} \\hMsun$, a factor of 100 above $M_{\\rm min}$), we do find a statistically acceptable fit, but even this model underpredicts the one-halo term of $\\wrp$ while overpredicting the two-halo term.  The difficulty in reproducing $\\wrp$ for faint red galaxies could signify a breakdown of our assumption that the HOD is independent of large-scale environment, but the obvious forms of environment dependence (redder central galaxies in older halos) go in the wrong direction. Clearly the clustering of the faint red galaxy population merits further study.  While we do not know of any planned redshift surveys that will provide better statistics for such low luminosity galaxies, cross-correlation of photometric samples with redshift samples of more luminous galaxies may allow higher precision clustering measurements from a larger effective volume.  Our measured luminosity and color trends agree with those found in earlier studies, most notably the \\cite{norberg01,norberg02} studies of the 2dFGRS and the Z05 study of SDSS DR2, but the SDSS DR7 sample allows measurements of higher precision, greater detail, and wider dynamic range. Our conclusions about the luminosity and color dependence of galaxy halo occupations are generally consistent with those found in earlier studies (e.g., \\citealt{bosch03a,collister05,yang05b}; Z05), although the greater precision and dynamic range of our clustering measurements allows us to examine this dependence in substantially greater detail. Even with SDSS DR7, the tests in \\S\\ref{subsec:cosmicvariance} reveal significant finite-volume effects for samples with limiting absolute magnitude $M_r \\approx -20$, which extend just far enough to enclose the Sloan Great Wall, and for samples with limiting magnitude $M_r \\geq -18$, which have small total volume.  These effects have a significant influence on the $(r_0,\\gamma)$ values of power-law fits to these samples (and their color-defined subsamples). They have little impact on the best-fit values of HOD parameters, though they do affect the $\\chi^2$ values of HOD fits. The finite-volume uncertainties limit the strength of our conclusions about the faint red galaxy population.  Our modeling in this paper derives HOD parameters for well-specified classes of galaxies defined by thresholds or bins in luminosity and divisions in color. The related formalism of conditional luminosity functions \\citep{yang03} seeks to provide a continuous description of the dependence of the galaxy luminosity function on halo mass. In a subsequent paper (Zheng et al., in preparation), we will present a generalization of this approach to luminosity-color distributions and apply it to our $\\wrp$ data, resulting in a comprehensive model that synthesizes the information from all of the measurements presented here.  Our HOD parameterization is flexible enough to describe the predictions of galaxy formation models accurately \\citep{zheng05}, and with an assumed cosmological model the $\\wrp$ measurements are themselves sufficient to provide tight constraints on HOD parameters. Studies of other real-space clustering measures, such as the multiplicity function of groups, the three-point correlation function, the topology of isodensity surfaces, and the void probability function can test the HOD models presented here, perhaps revealing breakdowns of this parameterization that would point to new aspects of galaxy formation physics.  Most interesting would be to find evidence for environmental variations of the HOD, as this would tie observable galaxy properties to features of halo formation history that correlate with large-scale environment at fixed halo mass. Conversely, limits on environmental variations (e.g., \\citealt{blanton07,tinker08}) limit the degree to which galaxy properties can be driven by quantities such as halo formation time or concentration.  Uncertainties in cosmological parameters within the range allowed by other data have little impact on our conclusions. The largest effect is that changes to $\\Omega_m$ or $\\sigma_8$ would shift the mass scale of the HOD \\citep{zheng07a}. The combination of $\\wrp$ constraints with dynamical measures that are sensitive to the halo mass scale allows novel constraints on these cosmological parameters.  Efforts in this direction are underway, using cluster mass-to-light ratios, galaxy-galaxy lensing, and redshift-space distortions. We expect these analyses, together with the cluster abundance analysis of \\cite{rozo10}, to yield tight independent constraints on $\\sigma_8$ and $\\Omega_m$ with several systematic cross-checks.  These constraints, based on the inferred amplitude of dark matter clustering, are complementary to those derived from the large-scale {\\it shape} of the galaxy power spectrum \\citep{reid10}, which can themselves be sharpened by using HOD modeling to account for the effects of scale-dependent galaxy bias \\citep{yoo09}. The combination of these constraints with those derived from CMB data, Type Ia supernovae, baryon acoustic oscillations, and other cosmological observables will allow stringent consistency tests of the $\\Lambda$CDM cosmological model, at the few-percent level.  Surveys of the next decade will extend many of these techniques to higher redshifts, but the SDSS maps of structure in the present-day universe still have much to teach us about galaxy formation and the physics of the cosmos.  \\smallskip"
    },
    "1005/1005.2139_arXiv.txt": {
        "abstract": "Magnetars have been suggested as the most promising site for the origin of observed soft gamma-ray repeaters (SGRs) and anomalous X-ray pulsars (AXPs). In this work we investigate the possibility that SGRs and AXPs might be observational evidence for a magnetic phase separation in magnetars. We study magnetic domain formation as a new mechanism for SGRs and AXPs in which magnetar-matter separates into two phases containing different flux densities. We identify the parameter space in matter density and magnetic field strength at which there is an instability for magnetic domain formation. We conclude that such instabilities will likely occur in the deep outer crust for the magnetic Baym, Pethick, and Sutherland (BPS) model and in the inner crust and core for magnetars described in  relativistic Hartree theory.  Moreover, we estimate that the energy released by the onset of this instability is comparable with the energy emitted by SGRs. ",
        "introduction": "Soft gamma-ray repeaters (SGRs) are compact objects undergoing episodic instabilities which produce  super-Eddington X-ray outbursts. Up to now $6$ SGRs (4 confirmed and 2 candidates) have been observed\\footnote{http://www.physics.mcgill.ca/$\\sim$pulsar/magnetar/main.html}. They are believed to be a new class of $\\gamma$-ray transients that are different from the source of ordinary gamma-ray bursts. Observations of the spin-down time-scale \\citep{kouveliotou1} have confirmed the fact that these SGRs are newly born neutron stars with a very large surface magnetic field ($B \\sim 10^{15}$ G). Such stars have been named magnetars \\citep{duncan, kouveliotou2}. About 10 AXPs are also categorized as magnetars \\citep{kaspi, paradijs}. Even though the magnetar model is generally accepted as the paradigm for both SGRs and AXPs, it is not easy to explain both objects simultaneously by a consistent set of parameters.  Woods et al.\\citet{woods} have reported evidence for a sudden magnetic field reconfiguration in SGR 1900+14 during the giant flare of August 27, 1998. This scenario requires a reorganization of the magnetic field both inside and outside the  star. Sharp field gradients are postulated to  create a fracture of the rigid outer crust of the neutron star. Cheng et al.\\citet{cheng} have shown that SGR events and earthquakes share four distinctive statistical properties: 1) power-law energy distributions; 2) log-symmetric waiting time distributions; 3) strong positive correlations between waiting times of successive events; and 4) weak or no correlation between intensities and waiting times. These statistical similarities, together with the fact that the crustal energy liberated by starquakes is sufficient in principle to fuel the soft gamma-ray flashes, suggest that SGRs are indeed powered by starquakes. Moreover, there is also a  strong correlation between the magnitude and  waiting times both for active earthquake regions and for SGR 1806-20. Thus, the statistical similarities between earthquakes and SGR events argue for physically similar origins. The Quasi-Periodic Oscillations (QPOs) observed at late-times of giant flares of SGR 1806-20, SGR 1900+14, and SGR 0525-66 \\citep{mareghetti, watts} constitute  another piece of observational evidence in favor of  starquakes. These QPOs are most likely due to seismic oscillations induced by the large crustal fractures occurring in extremely energetic events similar to what happens after earthquakes. Such oscillations could be limited to the crust or involve the entire neutron star.  Crackling noise \\citet{sethna} arises when a system responds to changing external conditions through discrete, impulse events spanning a broad range of sizes. Bak, Tang, and Wiesenfeld (1988) introduced a connection between dynamical critical phenomena and crackling noise.  They emphasized how systems may end up naturally at the critical point through a process of self-organized criticality. Based upon this idea of the crackling noise, Kondratyev  (2002) studied the statistics of magnetic noise in neutron star crusts, and compared its intensity and statistical properties to the burst activity of SGRs. He then argued that the noise could originate from magnetic avalanches. However, because of the required inhomogeneous crust structure, he postulated the existence of magnetic domains within the neutron-star crust for  an interior magnetic field strength in the range $\\sim 10^{16} - 10^{17}$ G. Using the randomly jumping interacting moment (RJIM) model, it was shown that the burst intensity and waiting time distributions are not only  in good agreement with observations, but also are analogous with the statistical properties of SGRs.  Whether or not  magnetars are the source of SGRs and AXPs, as relics of stellar interiors, the study of the magnetic fields in and around degenerate stars should give important information on the role such fields play in star formation and stellar evolution \\citep{suh01}. The scalar virial theorem implies an allowed internal field strength for a star of $B \\lesssim 2 \\times 10^8 (M/M_{\\odot}) (R/R_{\\odot})^{-2}$ G for a star of size $R$ and mass $M$. For a typical neutron star the maximum interior field strength could thus reach $B \\sim 10^{18}$ G. Since strong interior magnetic fields modify the nuclear equation of state for degenerate stars, their structure will also be changed \\citep{cardall}.  Even though SGRs appear to be observable consequences of starquakes or surface fractures, their detailed mechanism is not known. Moreover, it is unlikely that AXPs could only be explained by starquakes in strong magnetic fields. Therefore, in this paper we suggest a magnetic domain model  to correlate smoothly between the statistics of starquakes and magnetic avalanches in magnetar crusts. In this work magnetic properties of magnetar-matter such as the magnetization and the susceptibility are calculated in the framework of three different representative equations of state. We consider an ideal $n-p-e$ gas, relativistic Hartree mean field theory, and the magnetic Baym, Pethick, and Sutherland (BPS) model \\citep{lai}. It has been shown \\citep{broderick} that the magnetization of magnetar-matter undergoes large $de ~Haas-van ~Alphen$ oscillations. The magnetic susceptibility can then lead to a  region unstable to the formation of  {\\it magnetic domains}. It has not yet been demonstrated, however,  that magnetic domains actually form in magnetar-matter. Here we show that it is indeed possible to form such magnetic domains  in magnetars, and that  these could affect the surface properties and structure of magnetars  possibly leading to observable consequences such as starquakes, glitches,  and $X$- or $\\gm$-ray emission. ",
        "conclusions": "In this work we have studied magnetic domain formation as a new mechanism for SGRs and AXPs. In this paradigm magnetar-matter separates into two phases containing different flux densities. We have identified the parameter space in matter density and magnetic field strength at which there is an instability for magnetic domain formation and have shown  that such instabilities are  likely to occur in the deep outer crust for the magnetic BPS model, and in the deeper part of the inner crust and core for magnetars described by  relativistic Hartree theory. Moreover, we have estimated the strain on the outer crust induced by the formation of such domains and found that the anticipated energy release  is comparable with the energy emitted by typical SGRs. Hence, we propose that the magnetic domain formation scenario described here represents a new possible mechanism to drive the giant flares of SGRs as well as  X-ray outbursts and the quiescent phase of AXPs. At the very least,  this proposal warrants further investigation. Moreover, since the physical length variation caused by the magnetic domain formation might lead to solid crustal deformation and catastropic cracking, SGRs might be sources of gravitational waves (GWs) \\citep{abbott} even though there is not yet evidence of GWs associated with  observed SGR bursts. However, if it becomes possible to detect GWs from SGRs, that may be a way to verify the magnetic domain model in magnetars. Clearly, the next step is to undertake detailed dynamical numerical studies of the formation and evolution of such magnetic domains in neutron star crusts.  Efforts along this line are currently underway \\citep{suh03}."
    },
    "1005/1005.0620_arXiv.txt": {
        "abstract": "{Sun-like stars are promising candidates to host exoplanets and are often included in exoplanet surveys by radial velocity (RV) and direct imaging. In this paper we report on the detection of a stellar companion to the nearby solar-analogue star HD 104304, which previously was considered to host a planetary mass or brown dwarf companion.} {We searched for close stellar and substellar companions around extrasolar planet host stars with high angular resolution imaging to characterize planet formation environments. } {The detection of the stellar companion was achieved by high angular resolution measurements, using the ``Lucky Imaging'' technique at the ESO NTT 3.5m with the AstraLux Sur instrument. We combined the results with VLT/NACO archive data, where the companion could also be detected. The results were compared to precise RV measurements of HD 104304, obtained at the Lick and Keck observatories from 2001-2010. } {We confirmed common proper motion of the binary system.  A spectral type of M4V of the companion and a mass of 0.21\\,M$_\\odot$ was derived. Due to comparison of the data with RV measurements of the unconfirmed planet candidate listed in the Extrasolar Planets Encyclopaedia, we suggest that the discovered companion is the origin of the RV trend and that the inclination of the orbit of $i\\approx35\\degr$ explains the relatively small RV signal.} {} ",
        "introduction": "HD 104304 (LHS 5206, HIP 58576) is a high proper motion star at a distance of 12.91 $\\pm$ 0.13\\,pc \\citep{perryman97}. Because of its proximity, brightness, and observability from both hemispheres, it has been the subject of numerous studies. It is a chromospherically inactive star with log R$'$(HK) = $-4.933$ \\citep{gray06} and with an X-ray luminosity of L$_{\\rm X}$ = $1.2 \\times 10^{20}$\\,W \\citep{schmitt04}. It has a spectral type of G8IV to K0IV, log g = 4.35 to 4.15, and [Fe/H] = 0.22 to 0.18 \\citep{gray06,cenarro07}, and is a thin disk member \\citep{soubiran05}. With an age of $\\approx$8.5\\,Gyr and a mass of $\\approx$1.01\\,M$_\\odot$ \\citep{takeda07} it is considered a close analogue to the Sun at approximately twice its present age, and is a target of a high-precision Doppler survey aimed at identifying planetary companions to Sun-like stars \\citep{Wright2004}. HD 104304 has also been suggested as a target to search for habitable planets with future space missions \\citep{lopez05}.   In 2008 we started a high-angular resolution imaging survey of exoplanet hosts stars with the goal to identify close stellar and substellar companions. The first results of this survey were reported by \\citet{daemgen09}. HD 104304 was included in the survey because of the identification of a potential substellar companion by \\cite{Marcy2007}.  In the following we will use the designation HD 104304A for the primary star and HD 104304B for the companion. ",
        "conclusions": "The combination of NACO/SDI archival data and AstraLux Sur data made it possible to confirm the stellar companion around HD 104304. Based on the photometry, a spectral type of M4V and a mass of 0.21 M$_\\odot$ were derived for HD 104304B.  The primary star is well studied, e.g. it has a known metallicity of [Fe/H] = 0.22 to 0.18 \\citep{gray06,cenarro07}, an age of $\\approx$\\,8.5\\,Gyr \\citep{takeda07} and a mass nearly equal to that of our Sun. This knowledge is a valuable asset for the further characterisation of any stellar or substellar companion in this system.   By modeling the theoretical RV curve  and comparing it to the available RV observations obtained over a period of 9 years, we have shown that the observed RV trend could be explained by the detected stellar companion. A theoretical RV curve with an eccentricity of $e=0.29$, an orbital period of $P$ = 48.5\\,years and an inclination of the orbit of $i=35\\degr$~fits best the RV measurements. This gives evidence that the planet candidate listed in the Extrasolar Planets Encyclopaedia around HD 104304 is identical with the presented stellar companion. Astrometric monitoring of the binary system over the next few years is necessary to confirm this statement and to precisely determine the orbit and its orientation.  In the case of confirmation of a planet around HD 104304, this system would be highly interesting in the context of planet formation and migration around close binaries with solar analogue host stars."
    },
    "1005/1005.0599_arXiv.txt": {
        "abstract": "The newly developed X-ray visibility forward fitting technique is applied to Reuven Ramaty High Energy Solar Spectroscopic Imager (RHESSI) data of a limb flare to investigate the energy and height dependence on sizes, shapes, and position of hard X-ray chromospheric footpoint sources. This provides information about the electron transport and chromospheric density structure. The spatial distribution of two footpoint X-ray sources is analyzed using PIXON, Maximum Entropy Method, CLEAN and visibility forward fit algorithms at nonthermal energies from $\\sim 20$ to $\\sim 200$ keV. We report, for the first time, the vertical extents and widths of hard X-ray chromospheric sources measured as a function of energy for a limb event. Our observations suggest that both the vertical and horizontal sizes of footpoints are decreasing with energy. Higher energy emission originates progressively deeper in the chromosphere consistent with downward flare accelerated streaming electrons. The ellipticity of the footpoints grows with energy from $\\sim 0.5$ at $ \\sim 20$ keV to $\\sim 0.9$ at $\\sim 150$ keV. The positions of X-ray emission are in agreement with an exponential density profile of scale height $\\sim 150$~km. The characteristic size of the hard X-ray footpoint source along the limb is decreasing with energy suggesting a converging magnetic field in the footpoint. The vertical sizes of X-ray sources are inconsistent with simple collisional transport in a single density scale height but can be explained using a multi-threaded density structure in the chromosphere. ",
        "introduction": "In the standard flare scenario, electrons that were accelerated in the corona stream downwards toward the dense layers of the solar atmosphere, where they are stopped via collisions producing intense hard X-ray (HXR) emission in the chromosphere. Higher energy electrons penetrate deeper into the chromosphere. Therefore, measurements of the spatial structure of the HXR emission as a function of energy provide information about the chromospheric density. Being optically thin, X-rays give the most direct information about the spatial and energy distribution of energetic electrons in the solar atmosphere. Prior to the launch of the Ramaty High Energy Solar Spectroscopic Imager (RHESSI) \\citep{lin2002}, HXR instruments typically had limited imaging-spectroscopy capabilities \\citep{Kosugi1992yohkoh}. RHESSI's aptitude to image in various energy ranges opens new horizons for studying the detailed structure of HXR emitting sources in the chromosphere.  RHESSI does not directly image the Sun but uses 9 pairs of Rotating Modulation Collimators (RMCs) to time-modulate spatial information in the signal obtained in its germanium detectors \\citep{hurford2002}. Each RMC has a different thickness of its slits and slats making it sensitive to different spatial scales, providing modulation at nine spatial frequencies. The reconstruction of an image from these time modulated lightcurves, can be accomplished by various imaging algorithms \\cite[][]{emslie2003,BattagliaBenz2007,KruckerLin2008,Saint-Hilaire_etal2008,dennis2009}. The new visibility based approach to RHESSI imaging starts by summing (stacking) the lightcurves per roll bins over a few spin periods of the spacecraft \\citep{schmahl2007}. The fitted amplitudes and the phases in the individual roll bins are X-ray visibilities. This effectively provides two dimensional spatial Fourier components (X-ray visibilities) over a wide range of energies \\citep{hurford2002,schmahl2007}. To convert the time-modulated signal or X-ray visibilities to an image is an inverse problem \\citep[e.g.][]{Piana_etal2007,Prato_etal2009}. The reconstructed images face unavoidable difficulties due to measurement errors, finite coverage in Fourier space, and ill-posedness of the reconstruction problem. The resulting reconstruction errors and small dynamic range makes it difficult to accurately measure source sizes from reconstructed images. At best the imaging resolution is down to 2 arcseconds but in practice is around 7 arcseconds for the typical flare nonthermal energy range.  However, the moments of X-ray source distribution (source position, source size, etc) can be inferred with higher precision either from the time-modulated signal \\cite[][]{aschwanden2002,KruckerLin2008,Saint-Hilaire_etal2008,Fivian_etal2009} or visibilities \\citep{Xu_etal2008,Kontar_etal08,dennis2009,Prato_etal2009}. Thus, RHESSI measurements of X-ray source positions  can recover sub-arcsecond information using RHESSI modulated lightcurves \\citep{aschwanden2002,Liu_etal2006,Mrozek06} or visibilities \\citep{Kontar_etal08,dennis2009,Prato_etal2009}. This has allowed clear demonstration of the height-energy dependence of HXR sources: higher energy sources originate at lower heights \\citep{aschwanden2002,brown2002,Liu_etal2006}. This has substantially improved upon previous results with Yohkoh/HXT \\citep{1992PASJ...44L..89M}. The recently developed visibility-based technique \\citep{schmahl2007}, allowed \\citet{Kontar_etal08} to improve previous measurements and infer characteristic sizes (FWHM) of HXR footpoints and not just the centroid height. From this the convergence of the magnetic flux and neutral hydrogen density distribution in the chromosphere was inferred.  In this paper we study the structure of solar flare hard X-ray sources using various imaging algorithms: PIXON, MEM, CLEAN and visibility forward fit. Using RHESSI X-ray visibilities we find the characteristic shapes and positions for different energy ranges. We show that the technique of forward fitting X-ray visibilities allows us to determine not only the FWHM of the sources but vertical and horizontal sizes of the sources, which is required for examining the density structures of the chromosphere. Theoretical relationships were compared with observations to find the density structure of the chromosphere. The vertical size of the X-ray sources is found to be larger than the ones predicted by a hydrostatic atmosphere in thick-target scenario. However, assuming that the electrons are propagating along several narrow threads with different density profiles can explain the measured vertical sizes of the sources. ",
        "conclusions": "Using X-ray visibility forward fits, we inferred not only the characteristic sizes and positions but the shapes of HXR sources. The January 6th 2004 event indicates an overall decrease in size of the source and increase of the source ellipticity with energy. The FWHM of the southern source decreases from $\\sim6.5''$ down to around $\\sim 4''$ while the ellipticity of the source grows from $0.5$ up to $0.9$. The source is elongated along the limb as evident in nearly zero angle between semi major-axis and the limb, such orientation of the source is observed for all energy ranges. The vertical extent of the source is decreasing by a larger fraction (from $6''$ down to $3''$) than the horizontal size (from $7''$ down to $5.5''$) leading to larger elongation of the source along the limb. Hence the FWHM of the magnetic flux tube containing energetic electrons (semi-major axis) changes from $\\sim 9.7$~Mm at height $\\sim 1.1$~Mm down to $\\sim 7.6$~Mm at $0.6$~Mm above the photosphere. The northern footpoint is fainter, but shows a similar trend: the higher energies appear at low heights and the size of the source is decreasing with energy suggesting convergence of the magnetic field lines along which electrons propagate. Using X-ray visibilities we also re-analyzed a flare that occurred on February, 20th, 2002 that was previously studied by forward fitting time-modulated lightcurves \\citep{aschwanden2002}. Although this flare is rather weak, we found similar results: the higher energy sources appear at lower heights. The uncertainties are larger than in the January 6th, 2004 event but the density model proposed by \\citet{aschwanden2002} is within our error bars. We also note that the size of the sources in the February 20th flare decreases with energy similar to the event on January 6th.  Analysing HXR emission from footpoints we found that various imaging algorithms (PIXON, Visibility Forward Fit, CLEAN, MEM-NJIT) give generally similar results for spatial distributions of X-ray footpoints. Although RHESSI imaging algorithms can be adjusted by the parameter choice and hence X-ray images could be somewhat altered, there is a general trend for the algorithms. CLEAN with default set of parameters has a tendency to provide larger sources while MEM-NJIT tends to over-resolve X-ray sources. PIXON and visibility forward fit show very similar results. Visibility forward fit allows us to study sub-arcsecond distribution of hard X-ray sources in suitably orientated bright flares. Due to systematic differences in sizes we obtained for January 2004 and February, 2002 flares, we suggest that CLEAN and MEM-NJIT should be used for source size/shape measurements of X-ray sources with extreme caution.  The northern footpoint could be partially occulted as the highest energy photons come predominantly from the southern footpoint, but we note that this bright footpoint is unlikely to be occulted. As pointed out by G. Hurford\\footnote{Presentation at 9th RHESSI workshop in Genoa http://sprg.ssl.berkeley.edu/~krucker/genoa/position/XrayLimb-Genoa.ppt}, a source partially subtended by the solar disk should have a sharp edge, where the brightness of the source will drop from maximum to zero over rather small radial distance. The derivative of the source brightness in the $x$-direction (perpendicular to the limb), $\\partial I(x,y)/\\partial x$ will have a maximum at the limb, where $I(x,y)$ has a sharp drop. It is evident from Equation (\\ref{eq:vis}) that the corresponding visibilities, $2\\pi iuV(u,v)$, should have a well pronounced maximum, which should be evident in the measured amplitudes of visibilities. Specifically, the finer grid RMCs should show large amplitudes when the grids are parallel to limb, i.e. the visibility amplitudes are much larger at the phase angles $~0^o$ and $~180^o$, which is not evident in the event under study (cf Figure \\ref{fig:ff_fit}). To make the discussion more complete, we note that the line of sight effects cannot be definitively ruled out. The footpoints might be projections of two rather long flare ribbons viewed almost parallel to the line of sight, so that the vertical extension is the projection of different height. Finally we should note that if the occultation height is small, $\\lesssim0.5''$, our height measurements are lower limits, but the major conclusion about the vertical extend is the same. If, though unlikely, the lower part of the loop (footpoints) is occulted, these observations provide an interesting question to the flare models as to why the sources sizes decrease with energy and the higher energy sources appear lower and not at the lowest visible location where the density is the highest.  Our measurements show that while the locations of the maxima of X-ray emission are consistent with simple collisional transport in single density scale height chromosphere, the vertical sizes do not agree with the assumption of field aligned electron transport. The vertical extent of X-ray sources is 3-6 times larger than in the purely collisional model in single-density-scale-height chromosphere. However, a chromospheric model involving multiple density threads within the flux tube of a footpoint can explain both the position of the maximum and the vertical size of the sources. We note that pitch angle scattering due to Coulomb collisions is likely to be insufficient to produce so strong expansion. The X-ray source size increase due to collisional pitch angle scattering will be about a quarter of the electron stoping depth for initially field aligned electrons \\citep{Conway2000}. However, strong non-collisional scattering or wave-particle interactions \\citep[e.g.][]{Hannah_etal2009} might boost the vertical source sizes to the measured value and hence cannot be excluded and will also be consistent with lack of downward anisotropy found in X-ray flare emission \\citep{KontarBrown06mirror,Kasparova_etal2007}. We note that the adopted model does not account for the magnetic field and its effects on particle transport, which could lead to larger source sizes and is subject of additional modeling."
    },
    "1005/1005.5375_arXiv.txt": {
        "abstract": "We propose a new algorithm for solving a system of two nonlinear transcendental equations with two complex variables based on the M\\\"uller algorithm. The two-dimensional M\\\"uller algorithm is tested on systems of different type and is found to work comparably to Newton's method and Broyden's method in many cases. The new algorithm is particularly useful in systems featuring the Heun functions whose complexity may make the already known algorithms not efficient enough or not working at all. In those specific cases, the new algorithm gives distinctly better results than the  other two methods.  As an example for its application in physics, the new algorithm was used to find the quasi-normal modes (QNM) of Schwarzschild black hole described by the Regge-Wheeler equation. The numerical results obtained by our method are compared with the already published QNM frequencies and are found to coincide to a great extent with them. Also discussed are the QNM of the Kerr black hole, described by the Teukolsky Master equation. ",
        "introduction": "Solving  a system of two complex-valued nonlinear transcendental equations numerically is a task with varying difficulty, depending on the non-linearity of the system, the types of functions involved and the dimension of the space determined by the system. There are many well-known iterative root-finding algorithms, but most of them are specialized and optimized to work with a narrow set of functions -- for example polynomials or functions with real-valued roots. The two most heavily relied upon one-dimensional algorithms -- the secant method and Newton's method (or the Newton-Raphson method, \\cite{Newton, Raphson, Muller,numerical}) can work with a wide set of complex valued functions under proper conditions (see \\cite{numerical}), but they have their weak sides. Newton's method requires the evaluation of the function and its first derivative at each iteration. This increases the computational cost of the algorithm and it makes the algorithm unusable when the procedure evaluating the derivative of the function has numerical problems (for example see the discussion for the Heun functions below) or when derivative becomes zero or changes sign. The secant method avoids this limitation, but in the general case, it has lower order of convergence ($\\sim 1.618$) compared to that of Newton's method ($=2$) and the convergence of both of them is strongly dependent on the initial guess.  These problems of the algorithms are inherited by their multi-dimensional versions such as the generalized Newton-Raphson method (\\cite{Muller}) and the generalized secant method (Broyden's method, \\cite{Broyden}). Although those problems can have varying severity, there are systems in which those algorithms cannot be used effectively. There are also some novel approaches (see \\cite{multi}, \\cite{cgrasp}), but when they rely on the same one-dimensional algorithms, they are likely to share their weaknesses as well. Clearly there is a need for new algorithms that will enlarge the class of functions we are able to work with efficiently.  A great challenge in front of root-finding algorithms in modern physics can be found in systems including the Heun functions. The Heun functions are unique particular local solutions of a second-order linear ordinary differential equation from the Heun type \\cite{heun,heun1_,heun2_,heun3_} which in the general case have 4 regular singular points. Two or more of those regular singularities can coalesce into an irregular singularity leading to confluent differential equations and their solutions: confluent Heun function, biconfluent Heun function, double confluent Heun function and triconfluent Heun function. The Heun functions generalize the hypergeometric function (and also include the Lame function, Mathieu function and the spheroidal wave functions \\cite{heun2_,heun3_}) and their wide applications in physics (\\cite{heun3_}) was summarized recently in \\cite{heun1}. From that paper and the cited therein, it is clear that the Heun functions will be encountered more and more in modern physics from quantum mechanics to astrophysics, hence we need adequate numerical algorithms able to deal with them.  The work with the Heun functions, however, is more than complicated. While there are analytical works on the Heun functions, they were largely neglected until recently and therefore the theory is far from complete. The only software package currently able to work with them is \\textsc{maple}. Although, the routines that evaluate them in that package work well in the general case, there are some peculiarities -- there are values of the parameters where the routines break down leading to infinities or to numerical errors. The situation with the derivatives of the Heun  functions is even worst -- in some cases, they simply do not work or their precision is lower than that of the Heun function itself. Also, in some cases there are no convenient power-series representations and then the Heun functions are evaluated in \\textsc{maple} using numerical integration. Therefore the procedure goes slowly in the complex domain (compared to the hypergeometric function) which means that the convergence of the root-finding algorithm is essential.  Despite all the challenges in the numerical work with the Heun functions, they offer many opportunities to modern physics. For example, they occur in the problem of quasi-normal modes (QNM) of rotating and non-rotating black holes. In this case, one has to solve a two-dimensional connected spectral problem with two complex equations in each of which one encounters the confluent Heun functions. The analytical theory of the confluent Heun function is more developed than that of the other types of Heun functions, but still many unknowns remain. Again, the evaluation of the derivative of the confluent Heun function is problematic in \\textsc{maple} and the behavior of the function is hard to predict. In that situation Newton's method cannot be used as a root-finding algorithm. Broyden's algorithm works well in most cases, but it is slowly convergent even close to a root. It is clear,  then, that  we need a novel algorithm, that will offer quicker convergence than Broyden's algorithm, but without relying on derivatives.  To solve this problem, in the case of a system of two equations in two variables, our team developed a two-dimensional generalization of the M\\\"uller algorithm.  The one-dimensional M\\\"uller algorithm (\\cite{Muller2}) is a quadratic generalization of the secant method, that works well in the case of a complex function of one variable. It has very good convergence for a large class of functions ($\\sim 1.84$) and it is very efficient when the starting point (the initial guess) is close to a root. It is also well convergent when working with special transcendental functions. The two-dimensional M\\\"uller algorithm seems to inherit some of the advantages of its one-dimensional counterpart like good convergence and usability on large class of functions as our tests show. The new algorithm was used to solve the QNM problem in the case of a Schwarzschild black hole and it proved to work without significant deviations from the results published by Andersson (\\cite{Q_N_M}) and Fiziev (\\cite{Fiziev1}). Also, preliminary results for the QNM of the Kerr black hole are discussed and for them we also obtain a very good coincidence with published results \\cite{special31}.  The article is organized as follows: Section 2 reviews the one-dimensional M\\\"uller algorithm and its two-dimensional generalization, in Section 3 we discuss some physical application of the method and the numerical results obtained with it and in Section 4 we summarize our results. In the Appendix, the new algorithm is tested on various additional and more simple examples to verify its functionality. ",
        "conclusions": "We presented the general idea of a method for solving a system of two complex-valued nonlinear transcendental equations with complex roots based on the one-dimensional M\\\"uller method. As it was shown, the new method inherits some of the advantages of its one-dimensional counterpart as good convergence for a large class of functions. The efficiency of the new method is illustrated on different examples (For examples including elementary functions and simple special functions see Appendix.).  The numerical results showed that in many cases, it is comparable to Newton's method and Broyden's method, while in some specific systems  -- like those featuring Heun functions -- it is quicker than the other two. The complete mathematical investigation of the proposed new method, and especially its theoretical order of convergence under proper conditions on the class of functions $F_1,F_2$ is still an open problem.  Furthermore, the two-dimensional M\\\"uller algorithm works well with special functions like the confluent Heun function in terms of which one can solve analytically the Regge-Wheeler equation \\cite{Fiziev1}, the Zerilli  equation \\cite{Fiziev2}, the Teukolsky radial and angular equations \\cite{Fiziev3} and many others. Using this algorithm, we were able to solve directly the problem of quasi-normal modes of a Schwarzschild black hole with higher precision than that of the Broyden method. The so found solutions agree to great extent with previous published numerical results thus confirming the usefulness of the method. The new algorithm was used to study the stability of the so found QNM with respect to small variations in the phase condition in the radial variable and some preliminary results were presented. For other applications of the method see the recent references \\cite{PP, PP1}.  Future application of the two-dimensional M\\\"uller method would be in the case of QNMs of a Kerr black hole described by the radial and angular Teukolsky equations, which will be published elsewhere."
    },
    "1005/1005.3553_arXiv.txt": {
        "abstract": "We report the first simultaneous zJHK spectroscopy on the archetypical Seyfert~2 Galaxy NGC~1068 covering the wavelength region 0.9 to 2.4$\\mu$m. The slit, aligned in the NS direction and centred in the optical nucleus, maps a region 300~pc in radius at sub-arcsec resolution, with a spectral resolving power of 360~km~s$^{-1}$. This configuration allow us to study the physical properties of the nuclear gas including that of the north side of the ionization cone, map the strong excess of continuum emission in the K-band and attributed to dust and study the variations, both in flux and profile, in the emission lines. Our results show that (1) Mid- to low-ionization emission lines are splitted into two components, whose relative strengths vary with the position along the slit and seem to be correlated with the jet. (2) The coronal lines are single-peaked and are detected only in the central few hundred of parsecs from the nucleus. (3) The absorption lines indicate the presence of intermediate age stellar population, which might be a signicant contributor to the continuum in the NIR spectra. (4) Through some simple photoionization models we find photoionization as the main mechanism powering the emitting gas. (5) Calculations using stellar features point to a mass concentration inside the 100 - 200 pc of about 10$^{10}~M_{\\odot}$. ",
        "introduction": "Near-IR spectroscopy has been playing a unique role in our understanding of the AGN phenomenon for several reasons. First, it includes a wealth of emission lines and stellar absorption features not observed in the optical, which are useful for studying the physical properties of the emission gas and to determine the age and metallicity of the nuclear stellar population as, for instance, the strong CO bandheads in the H and K bands (e.g., Schreiber 1998; Origlia, Moorwood, \\& Oliva 1993). Second, extinction by dust is attenuated by a factor of ten relative to that of the optical, allowing to probe depths unreachable at shorter wavelengths. Third, it is a transition region, where the continuum from the central source no longer dominates while the thermal continuum produced by dust and stars becomes important.  In the last years, the interest for observation in this interval has increased, and thanks to the availability of cross-dispersed (XD) spectrographs, it is possible to study the whole 0.8-2.4$\\mu$m region in a single observation, avoiding the aperture and seeing effects that usually affects $JHK$ spectroscopy done in long-slit single band observations.  With the above in  mind, here we present the first spatially resolved XD spectroscopy covering the interval 0.8-2.4$\\mu$m made for NGC1068. This object is one of the nearest and probably the most intensely studied Seyfert 2 galaxy. Observations in all wavelength bands from radio to hard X-rays have formed a uniquely detailed picture of this source. NGC 1068 has played a unique role in the development of unified scenarios for Seyfert 1 and Seyfert 2 galaxies (Antonucci \\& Miller 1985), in the study of molecular gas in the nuclear region of Seyferts (e.g., Myers \\& Scoville 1987; Tacconi et al. 1994), and in elucidating the importance of star formation activity coexistent with the active galactic nucleus (AGN), on both larger (e.g., Telesco \\& Decher 1988) and smaller (Macchetto et al. 1994; Thatte et al. 1997) scales. NGC 1068 also hosts a prominent narrow-line region (NLR) that is approximately co-spatial with a linear radio source with two lobes (Wilson \\& Ulvestad 1983). The narrow emission line region has been extensively characterized from subarcsecond clouds probed by the Hubble Space Telescope  (HST) (Evans et al. 1991; Macchetto et al. 1994) to the ionization cone and extended emission-line region (Pogge 1988; Unger et al. 1992) extending to radii of at least 30$\\arcsec$  (1$\\arcsec$=72 pc at the distance  of 14.4 Mpc; Tully 1988).  From approximately 2\" southwest to the nucleus to 4\" northeast, emission lines exhibit multiple components (Cecil et al. 1990; Crenshaw \\& Kraemer 2000a). Broad lines were found to be approximately 2500 - 4000 km s$^{-1}$ wide, which may be linked to those that are  found in polarized light and believed to be reflected light from the  inner broad line region (BLR). Narrow lines consist of a pair of red and blue components. Optical studies of [O\\,{\\sc ii}] and [N\\,{\\sc ii}] line profiles suggest that the  separation of these two components varies across the conical NLR.  Between $\\sim$2\".5 and 4\".5 northeast from the nucleus, UV line emission is redshifted relative to the systemic value, a pattern that is interpreted as the expansion of the plasma in the radio lobe (Axon et al. 1998).  The mechanisms powering the gas of the NLR, i.e. photoionization from the nucleus or shocks produced by jets, have long been under debate. Dopita \\& Sutherland (1996) and Bicknell et al. (1998) proposed that the emission in the NLR may be entirely caused by shocks. Velocity splitting over 1000 km$^{-1}$, reported by Axon et al. (1998), in the vicinity of some of the bright emission-line knots provides evidence  that fast shocks exist in the NLR of NGC 1068. However, more recent HST  data (Crenshaw \\& Kraemer 2000b; Cecil et al. 2002; Mazzalay et al. 2009) found that the emission line ratios are consistent with photoionization instead of shock heating mechanism. On the other hand, models accounting for  both photoionization from the central radiation and shocks were required to explain both the continuum and emission-lines observed in the spectra of active galaxies (e.g., Contini, Rodr\\'{\\i}guez-Ardila \\& Viegas 2003, Rodr\\'{\\i}guez-Ardila, Contini  \\& Viegas 2005).   The scenario in NGC~1068 is thus clearly complex, and observational constraints from all wavelengths need to be put together if we want to understand the processes taking place in this galaxy. Spatially resolved NLR spectroscopy in the NIR can contribute to the solution of the NGC~1068 puzzle. Here, we present the first observations on this source covering simultaneously the 0.8$-2.4~\\mu$m region at moderate spectral and spatial resolution. Our data are thus able to map a wavelength interval rich in emission lines with a large range of ionization and absorption features from the stellar population not only from the nucleus but also from the circumnuclear region, poorly studied in the literature. In \\S 2 a description of the observations is given. The main characteristics of the nuclear and  extended spectra are presented in \\S 3 and \\S 4 , respectively.  Photoionization model predictions for the most intense lines are discussed in \\S 5, whereas  the calculations of the central mass based on stellar features appear in \\S 6. Finally, our concluding remarks are presented in \\S 7. ",
        "conclusions": "We carried out long-slit spectroscopy in the Near-IR of NGC~1068 and presented here for the first time the whole range from 0.9 to 2.4$\\micron$, extending 15\" in the nuclear region. The nuclear emission shows many important emission lines, from low ionization to coronal lines, molecular lines and absorption features. We show that many emission lines are composed by two narrow components, one in the galaxy rest frame, and the other blueshifted in the south and redshifted in the north apertures. The relative intensity of the peaks also varies and they seem to be correlated with the jet of the galaxy. No broad component ($\\sim$~2000) was identified in any line profile. Not all emission lines have double peaked profiles and some lines are present only near the center, particularly the high ionization lines. This is a clear indication that the emitting gas is more ionized in the nuclear region.  We estimated the extinction based on Paschen hydrogen and [FeII] emission lines. E(B-V) is highest in the nucleus then in the adjacent regions, but there are peaks around 150 - 200 pc from each side of the nucleus. This is also close to the region where the second components of the emission lines have a maximum, and may be associated with dense clouds around these regions. H$_2$ emission is very extended, but it is difficult to identify the ionization process because only the 1-0(S1) line could be measured.  Simple photoionization models were used to try to explain the emission lines, and a first approximation shows that the line ratios can be explained with pure photoionization, as long as low density clouds are present. A more detailed analysis however is necessary, to rule out other ionization mechanisms.  We also used the stellar signatures to determine the mass of the central region, obtaining a stellar core of 1.5$\\times$10$^{10}$M$_{\\sun}$."
    },
    "1005/1005.3848_arXiv.txt": {
        "abstract": "In this work we study the effects of field space curvature on scalar field perturbations around an arbitrary background field trajectory evolving in time. Non-trivial imprints of the `heavy' directions on the low energy dynamics arise when the vacuum manifold of the potential does not coincide with the span of geodesics defined by the sigma model metric of the full theory. When the kinetic energy is small compared to the potential energy, the field traverses a curve close to the vacuum manifold of the potential. The curvature of the path followed by the fields can still have a profound influence on the perturbations as modes parallel to the trajectory mix with those normal to it if the trajectory turns sharply enough. We analyze the dynamical mixing between these non-decoupled degrees of freedom and deduce its non-trivial contribution to the low energy effective theory for the light modes. We also discuss the consequences of this mixing for various scenarios where multiple scalar fields play a vital role, such as inflation and low-energy compactifications of string theory. ",
        "introduction": "\\setcounter{equation}{0} A thorough and tractable understanding of early universe physics through an ultraviolet (UV) complete description, such as string theory, remains out of our reach for now. While we do not completely understand the precise features of the UV completion of the standard model coupled to gravity, or even the gross structure of such a putative theory, one of its likely generic consequences is the presence of a large number of scalar fields and possibly a large number of vacua. The masses of these fields are typically associated with the internal structure of the theory, for example, the cutoff scale of the effective action which encapsulates the UV relevant physics that completes our theory.  However, in general it might be that different fields have different mass scales associated with them such that a hierarchy appears in the field content, and we can colloquially speak of `light' and `heavy' modes. Usually, the heavy fields are assumed to be integrated out since they are kinematically inaccessible provided that their masses are larger than the scale of physics of interest. However in reality, we cannot completely ignore the heavy fields as the conditions underlying the decoupling theorem of Ref.~\\cite{Appelquist:1974tg} can sometimes be relaxed, for instance through time dependence in the heavy sector, or through dynamical mixing of heavy and light sectors. In the context of field theories at a point in field space, gauge symmetry breaking can lead to a non-decoupling scenario (see~\\cite{SekharChivukula:2007gi}, and references therein, for a recent example). In the context of supergravity, decoupling of heavy fields is still actively being studied~\\cite{Choi:2004sx,deAlwis:2005tg,deAlwis:2005tf,Binetruy:2004hh,Achucarro:2007qa,Achucarro:2008sy,Achucarro:2008fk,BenDayan:2008dv,Gallego:2008qi, Brizi:2009nn, Gallego:2009px, Gallego:2011jm}.  \\begin{figure} \\includegraphics[width=0.48\\textwidth]{images/ellipspot.png} \\caption{A curved trajectory with a sudden bend, where a scalar potential presents a turning flat direction with a massive mode remaining perpendicular to it} \\label{fig:curved_traj} \\end{figure} In a series of papers, this collaboration will examine the conditions under which decoupling is relaxed, or fails outright, and explore quantitatively the consequences of this failure.  The subject of this paper is to study multi-scalar field theories in Minkowski spacetimes characterized by having a hierarchy of scales between different families of fields. Our main goal is to characterize the interaction between light and heavy degrees of freedom as the background expectation value of the scalar fields (which we assume to be spatially homogeneous) varies in time, traversing a curved trajectory in field space. Specifically, we analyze those situations where the heavy degrees of freedom correspond to the transverse modes with respect to the path followed by the background fields. Pictorially, these curved trajectories are placed in the nearly flat valleys of the vacuum manifold. An example of such a curved trajectory is schematically shown in Fig.~\\ref{fig:curved_traj} where a scalar potential presents a curved flat direction with a massive mode remaining perpendicular to it. Such curvature mixes light and heavy directions along the field trajectory, possibly non-adiabatically if the field space is `curved' enough. In the context of inflation, the effects studied in this work lead to potentially observable effects in the spectrum of primordial perturbations (see \\cite{Achucarro:2010da} and references therein).     In the case of two-field models, we will show that whenever the trajectory makes a turn in field space such that the light direction is interchanged by a heavy direction (and vice versa), the low energy effective action describing the dynamics of the light degrees of freedom generically picks up a non-trivial correction, hereby parameterized by $\\beta$ and given by \\be \\beta = \\ln \\left( 1+ \\frac{4\\dot \\phi_0^2}{ \\kappa^2 M_H^2} \\right) \\, , \\ee where $\\dot \\phi_0$ is the velocity of the background scalar field, $\\kappa$ is the radius of curvature characterizing the bending of the trajectory in field space,  and $M_H$ is the characteristic mass scale of the heavy modes.  We emphasize that the bending is measured with respect to the geodesics of the sigma model metric, whether or not it is canonical; notice that $\\kappa$ represents a distance measured in the geometrical space of scalar fields, and therefore it has units of mass.  More precisely, at low energies ($k \\ll M_H$) the equation of motion describing the scalar field fluctuation {\\em parallel} to the trajectory of background fields is found to be \\be \\label{effective-eq} \\ddot \\varphi - e^{-\\beta} \\nabla^2 \\varphi + M^2_L \\, \\varphi = 0 \\, , \\ee with $\\nabla^2 \\equiv \\delta^{ij}\\partial_i\\partial_j$ being the spatial Laplacian: see (\\ref{low-energy-eff-action}). Thus the kinetic energy term is modified by a factor of $e^{-\\beta}$ in front of the spatial gradient term, in agreement with \\cite{Tolley:2009fg}. In the context of inflation, this can give rise to sizable effects that may be observable in the next generation of the cosmic microwave background observations. For example, for inflation in supergravity models, the masses of heavy degrees of freedom during inflation are typically of order $M_H = \\mathcal{O}(H)$, with $H$ being the Hubble parameter during inflation. This leads to the naive expectation \\be e^\\beta \\sim 1+ \\epsilon  \\frac{M_{\\rm Pl}^{2}}{ \\kappa^2} \\, , \\ee where $\\epsilon \\equiv -\\dot{H}/H^2 = (\\dot\\phi_0/H)^2/(2\\mpl^2)$ is the so-called slow-roll parameter. Then, if the bending of the trajectory is such that the radius of curvature $\\kappa$ becomes much smaller than $M_{\\rm Pl}$ one then is led to $e^\\beta \\gg 1$. Such a situation implies strong modifications on the scale dependence of inflaton fluctuations, contributing non-trivial features to the power spectrum well within the threshold of detection. The full analysis valid for inflationary cosmology is, however, more involved than the naive arguments above as it requires a detailed account of gravity and careful treatments of the multi-field dynamics at horizon crossing. For this reason, in the present note we study the basic features of the effects of the heavy-light field interactions during turns of the scalar field trajectory in a Minkowski background, and leave the study of inflationary de Sitter case for a separate paper~\\cite{Achucarro:2010da}. We stress here that the perspective that follows from our analysis is that the only relevant criteria for the effects we uncover (such as the modified speed of sound for the perturbations) is that our system deviates sufficiently from a geodesic trajectory \\footnote{In \\cite{Tolley:2009fg} \\cite{sera}, related effects were uncovered that were specifically induced by non-canonical kinetic terms. Our perspective is that the latter is more readily thought of as having been induced by non-geodesic trajectories in field space. Certainly any metric on field space can be brought into canonical form by a suitable field redefinition at the expense of introducing complicated derivative interactions at higher orders. Therefore the criteria must be more general than simply considering non-canonical kinetic terms (to induce for example, a reduced speed of sound for the scalar field perturbations).}.    Our formalism, and its application to inflation \\cite{Achucarro:2010da} builds upon previous work \\cite{GrootNibbelink:2000vx,GrootNibbelink:2001qt,Gordon:2000hv,Nilles:2001fg} that we extend to be able to account for the effects coming from very sharp turns and also prolonged turns in which the true ground state deviates significantly from its single-field approximation. By disregarding the effects of gravity, we are able to describe in detail the evolution of the quantum modes in some idealized settings (e.g. turns at a constant rate).  We also extend recent results by \\cite{Chen:2009we, Tolley:2009fg}, that we recover in the right limits. In particular, we are able to clarify the relevance or otherwise of non-canonical kinetic terms, in particular with respect to a reduced speed of sound in the effective low energy theory (see also \\cite{sera}). We expect our results to be relevant  for any phenomena where the variation of vacuum expectation values plays a significant role, such us inflation, phase transitions, soliton interactions and soliton-radiation scattering.   In the remainder of this paper, we will first elaborate on the general framework and study the background equations of motion in Section~\\ref{sec:background}. In Section~\\ref{pert-quant}, we will study the theory of perturbations and the quantization around a general field trajectory. In Section~\\ref{Hierarchies}, we will apply the theory to the case of a large hierarchy between light and heavy degrees of freedom that span a curved field space trajectory. To conclude, we will discuss some applications in Section~\\ref{conclusions}, in particular regarding inflation and the decoupling of heavy modes in supergravity. Some calculational details are provided in the appendices. ",
        "conclusions": "\\setcounter{equation}{0} \\label{conclusions} In this note, we considered the structure of low energy description of scalar field theories with a pronounced hie\\-rar\\-chy of mass scales. First, we set up a framework for describing a light field moving along a multi-field trajectory in field space. From this, we determined the background equations of motion, around which we can study perturbations. Finally, we deduced the effective theory describing light perturbations for the case in which the background field is following a curved trajectory in field space.  The main manifestation of the non-trivial mixing of the heavy and the light directions is in the appearance of the coefficient $e^{-\\beta}$ in front of the term $( \\nabla \\varphi)^2$ containing spatial derivatives in the action (\\ref{low-energy-eff-action}). First, since $\\beta \\ge 0$, the net effect of the bending of the background trajectory is to reduce the energy per scalar field quantum. This is due to the fact that during bending the light modes momentarily start exciting heavy modes, therefore transferring energy to them. Second, since $\\beta = \\ln ( 1 + 4\\zeta^2 / M^2)$, there are two effects competing against each other in this process. On one hand one has $\\zeta = V_N / \\dot \\phi_0$, or alternatively $ |\\zeta| = \\dot \\phi / \\kappa$, which may be interpreted as the angular speed of the background field along the curved trajectory. On the other hand, there is the mass of the heavy mode $M$, which must be excited by the light modes during the bending.   We stress that the magnitude of the effect is parametrized entirely by the parameter $\\beta$, which depends on the velocity of the background trajectory and the induced curvature along the scalar field trajectory through the field manifold, parametrized by $\\kappa$. Whether the kinetic terms are canonical or not is immaterial as this is merely a question of the chosen basis in field space.  In what follows we discuss two direct applications of our results: inflation and the non-decoupling of light and heavy modes in Supergravity.    \\subsection{Inflation} Understanding in detail how light and heavy modes remain coupled under more general circumstances could be particularly significant for cosmic inflation~\\cite{Guth:1980zm,Albrecht:1982wi,Linde:1981mu}. Indeed, although current observations are consistent with the simplest model of single field inflation, it is rather hard to conceive a realistic model where the inflaton field alone is completely decoupled from UV degrees of freedom. One way of addressing this issue is by studying multi-field scenarios where many scalar fields have the chance to participate in the inflationary dynamics~\\cite{Starobinsky:1986fxa}, despite of different mass scales among the inflaton candidates. Hence, there could exist certain phenomena related to inflation in which the effects studied in this report can be relevant. This has also recently been considered in Refs.~\\cite{Tolley:2009fg, sera, Chen:2009zp}. A more detailed discussion can be found in \\cite{Achucarro:2010da}.  First, note that the equation of motion deduced from the action (\\ref{low-energy-eff-action}) is given by \\be \\ddot \\varphi + e^{-\\beta} k^2 \\varphi + M^2_L \\, \\varphi = 0 \\, . \\label{effective-eq-2} \\ee For definiteness, let us focus on phenomena characterized by $\\left| \\dot \\beta \\right| \\ll k$ and consider the case in which the potential is flat enough so that $M^2_L \\ll k^2$ is satisfied. Then, the time variation of $\\beta$ along the trajectory is small enough to allow us to write the mode solution as \\be \\varphi (k,t) = \\frac{e^{\\beta/4}}{\\sqrt{k}} \\exp \\left[ i e^{-\\beta/2} k \\, t \\right] \\, , \\label{sol-beta} \\ee where the factor $e^{\\beta/4}/\\sqrt{k}$ is necessary in order to satisfy the commutation relation $[ \\varphi, \\dot \\varphi ] = i$. This factor coincides with the one found in (\\ref{expansion-v-1}) for the amplitude of light modes in the case where $\\beta$ is a constant. To continue, from (\\ref{sol-beta}) we can see that in the vacuum, the two point correlation function of the perturbation $\\varphi ({\\bf x}, t)$, has the form $\\langle \\varphi ({\\bf x}, t) \\varphi ({\\bf y},t) \\rangle \\propto e^{\\beta(t)/2}$. One direct consequence of this result is for inflation, where the amplitudes of scalar fluctuations freeze after crossing the horizon, i.e. when the physical wavelength $k^{-1}$ satisfies the condition $e^{- \\beta/2} k = H$. More precisely, if we generalize (\\ref{effective-eq-2}) to include gravity, we would conclude that the speed of sound of adiabatic perturbations is given by \\be c_s^2 = e^{- \\beta} \\, . \\ee Such an effect is known to produce sizable levels of non-Gaussianities~\\cite{Alishahiha:2004eh}. Additionally it modifies the power spectrum as \\be P(k) \\simeq e^{\\beta(k) / 2} P_*(k) \\, ,  \\label{power-beta-2} \\ee where $P_*(k)$ is the conventional power spectrum $P_*(k) \\propto k^{n_s -1}$ deduced in single field slow-roll inflation, and $\\beta(k)$ is the value of $\\beta(t)$ at the time $t$ when the mode $k$ crosses the horizon~\\footnote{In the constant curvature case $\\kappa=\\textrm{const.}$, $\\beta(\\kappa)$ is constant as well and we find an overall modulation of the power spectrum compatible with Ref.~\\cite{Chen:2009zp}.}. Certainly the more interesting case corresponds to a varying $\\beta$, and since $\\beta$ can be large, such an effect may be sizable and observable in the near future. In many scalar field theories, such as supergravity, the masses of heavy degrees of freedom during inflation are typically of order $M \\sim H$, leading to the relation \\be e^\\beta \\sim 1+ 4 \\epsilon  \\frac{M_{\\rm Pl}^{2}}{ \\kappa^2} \\, . \\ee If the bending is such that $\\kappa \\ll M_{\\rm Pl}$ (a feasible situation) one then could obtains effects as large as $\\beta \\sim 1$. In the case where a turn of the trajectory happens during a few $e$-folds, one then should be able to observe features in the power spectrum of $\\mathcal{O}(1)$, particularly by modifying the running of the spectral index as $d n_s / d \\ln k$, which otherwise would be of $\\mathcal{O}(\\epsilon^2)$. A detailed computation of this effect can be found in ~\\cite{Achucarro:2010da}, where the interaction between light an heavy modes during horizon crossing is examined more closely (see also \\cite{sera}).    \\subsection{Decoupling of light and heavy modes in supergravity} Our results can be also used to assess when a low energy effective theory, deduced from a multi-scalar field theory containing both heavy directions and light directions, is accurate enough. As discussed in full detail in Appendix \\ref{appendix-zero-order}, whenever the background fields are evolving (as in inflation) the only way of having a vanishing $\\beta$-parameter is for a trajectory to correspond to a geodesic in the full scalar field manifold $\\mathcal M$. It is clear that the only way of achieving this is by having some property relating the shape of the potential $V(\\phi)$ with the geometry of $\\mathcal{M}$. In the particular case of supergravity such a property is known to exist, and therefore one should expect supergravity theories rendering low energy effective theories for which $\\beta$ vanishes exactly. To be more precise, in $N=1$ supergravity the scalar field potential is given by \\be V = e^{G} \\left(\\gamma^{a \\bar b} G_a G_{\\bar b} - 3\\right) \\ , \\ee where $\\gamma^{a \\bar b}$ is the inverse of the K\\\"ahler metric $\\gamma_{a \\bar b} = \\partial_{a} \\partial_{\\bar b} G$ deduced out of the generalized K\\\"ahler potential $G$, which is a real function of the chiral scalar fields $\\phi^a$, $G_a = \\partial_{a} G$ and $G_{\\bar b} = \\partial_{\\bar b} G$. Now, consider a supergravity theory in which a set of massive chiral fields $\\phi^H$ satisfy the condition \\be G_H = 0 \\ , \\ee along a given hypersurface $\\mathcal S$ in $\\mathcal{M}$ parametrized \\emph{only} by the light fields $L$. This means that a surface $S$ is defined by $f(H,\\bar{H})=0$ rather that by a function $f(H,\\bar{H},L,\\bar{L})=0$~\\cite{deAlwis:2005tf} (see also~\\cite{Achucarro:2008sy}). Then it is possible to verify that the scalar fluctuations $\\phi^L$ parallel to $\\mathcal S$ are decoupled from the fields $\\phi^H$, rendering $\\beta = 0$. To appreciate this, observe first that at any point on the surface $\\mathcal S$ the K\\\"ahler metric $\\gamma_{a \\bar b}$ is diagonal between the two sectors. Indeed, since $G_H = 0$ holds at any point in the surface, then it must be independent of arbitrary displacements $\\delta \\phi^L$ along $\\mathcal S$. This implies that \\be \\partial_{\\bar L} G_H = \\gamma_{H \\bar L}  = 0  \\qquad {\\rm on} \\,\\, \\mathcal S. \\ee This condition automatically ensures that in the absence of a scalar field potential, the trajectory along $\\mathcal S$ will be a geodesic. It remains then to verify that the potential does not imply quadratic couplings between both sectors, and therefore leaves these geodesic trajectories unmodified. Consider the first and second derivatives of the potential \\begin{widetext} \\bea \\label{der-1} \\nabla_a V &=&  e^{G} \\Big(G_a + G^b \\nabla_a G_b \\Big) + V G_a \\, , \\\\ \\label{der-2} \\nabla_a \\nabla_{\\bar b} V &=&  e^{G} \\Big(\\gamma_{a \\bar b} + \\! \\nabla_a G_c \\nabla_{\\bar b} G^c \\! - \\mathcal{R}_{a \\bar b c \\bar d} G^c G^{\\bar d}   \\Big) \\! +   G_a V_{\\bar b}   +  G_{\\bar b} V_a  +   ( \\gamma_{a \\bar b} - G_a G_{\\bar b} )  V \\, ,  \\quad  \\\\ \\label{der-3} \\nabla_a \\nabla_b V &=& e^{G} \\Big( 2 \\nabla_{a} G_{b} + G^c \\nabla_a \\nabla_b G_c \\Big) + G_a V_{b} + G_{b} V_a + (\\nabla_{a} G_{b} - G_a G_{b} ) V \\, .  \\eea \\end{widetext} Since $\\gamma_{H \\bar L} = 0$ and $G_H=0$, from (\\ref{der-1}) one immediately obtains $V_H = 0$ in $\\mathcal S$. It is not difficult to notice that also $\\nabla_L V_H = \\nabla_{\\bar L} V_H = \\nabla_{L} V_{\\bar H} = \\nabla_{\\bar L} V_{\\bar H} = 0$ which also hinge on $\\gamma_{H \\bar L} = 0$ and $G_H=0$. All of these results put together imply that the heavy sector will not affect the light sector as long as the background trajectory remains on the geodesically generated surface $\\mathcal S$. Conversely, deviations from the condition $G_{H} = 0$, or a surface $S$ with $H \\neq \\mathrm{const.}$, will produce interactions leading to the appearance of the coupling $\\beta$ studied in the present work. A particularly interesting example is Ref.~\\cite{Gallego:2008qi}, where $O(\\epsilon)$ couplings between heavy and light fields in the superpotential result in suppressed, $O(\\epsilon^{2})$ terms in the effective action for the light fields. This result was obtained by expanding about a particular $H=\\mathrm{const.}$ configuration which, for constant light background fields, only deviates at $O(\\epsilon)$ from the true solution to the equations of motion. But along an arbitrary background $L(t)$ the deviation will exceed $O(\\epsilon)$ for displacements $\\Delta L / \\kappa > \\epsilon$ (due to the $\\Gamma^{H}_{LL} \\dot{L}^{2}$ term in the $H$ equation of motion), and the corrections to the effective action discussed in this paper become dominant."
    },
    "1005/1005.0977_arXiv.txt": {
        "abstract": "Astrophysical sources emit gravitational waves in a large variety of processes occurred since the beginning of star and galaxy  formation. These waves permeate our high redshift Universe, and form a background which is the result of the superposition of different components, each associated to a specific astrophysical process. Each component has different spectral properties and features that it is important to investigate in view of a possible, future detection. In this contribution, we will review recent theoretical predictions for backgrounds produced by extragalactic sources and discuss their detectability with current and future gravitational wave observatories. ",
        "introduction": "A stochastic background of gravitational waves (GWs) is expected to arise as a super-position of a large number of unresolved sources, from different directions in the sky, and with different polarizations. It is usually described in terms of the GW spectrum, \\[ \\Omega_{\\rm GW} = \\frac{f}{\\rho_{\\rm cr}}\\frac{d\\rho_{\\rm GW}}{df}, \\] \\noindent where $d\\rho_{\\rm GW}$ is the energy density of GWs in the frequency range $f$ to $f+df$ and $\\rho_{\\rm cr}$ is the critical energy density of the Universe.  Stochastic backgrounds of GWs are interesting to investigate because they (i) represent a foreground signal for the detection of primordial backgrounds (see the contribution of M. Maggiore in this conference), and (ii) provide important information on GW sources and on the properties of distant stellar populations.  The starting point in any calculation of the GW backgrounds produced by extragalactic sources is the redshift evolution of the comoving star formation rate density. This is generally inferred from observations which have been recently extended out to redshifts $z \\sim 10$ following the analysis of the deepest near-infrared image ever obtained on the Hubble Ultra Deep Field using the new/IR camera on HST (Bouwens et al. 2010). The detection of galaxies at $z = 8.4-10$, just 500-600 Myr after recombination and in the heart of the reionization epoch, provides clues about galaxy build-up at a time that must have been as little as 200-300 Myr since the first substantial star formation (at $z \\sim 15-20$). The existence of galaxies at $z \\sim 10$ shows that galaxy formation was well underway at this very early epoch. ",
        "conclusions": ""
    },
    "1005/1005.5249_arXiv.txt": {
        "abstract": "{It is thought likely that vast numbers of nanoflares are responsible for the corona having a temperature of millions of degrees. Current observational technologies lack the resolving power to confirm the nanoflare hypothesis. An alternative approach is to construct a magnetohydrodynamic coronal loop model that has the ability to predict nanoflare energy distributions.} {This paper presents the initial results generated by a coronal loop model that flares whenever it becomes unstable to an ideal MHD kink mode. A feature of the model is that it predicts heating events with a range of sizes, depending on where the instability threshold for linear kink modes is encountered. The aims are to calculate the distribution of event energies and to investigate whether kink instability can be predicted from a single parameter.} {The loop is represented as a straight line-tied cylinder. The twisting caused by random photospheric motions is captured by two parameters, representing the ratio of current density to field strength for specific regions of the loop. Instability onset is mapped as a closed boundary in the 2D parameter space. Dissipation of the loop's magnetic energy begins during the nonlinear stage of the instability, which develops as a consequence of current sheet reconnection. After flaring, the loop evolves to the state of lowest energy where, in accordance with relaxation theory, the ratio of current to field is constant throughout the loop and helicity is conserved.} {There exists substantial variation in the radial magnetic twist profiles for the loop states along the instability threshold. These results suggest that instability cannot be predicted by any simple twist-derived property reaching a critical value. The model is applied such that the loop undergoes repeated episodes of instability followed by energy-releasing relaxation. Hence, an energy distribution of the nanoflares produced is collated. This paper also presents the calculated relaxation states and energy releases for all instability threshold points.} {The final energy distribution features two nanoflare populations that follow different power laws. The power law index for the higher energy population is more than sufficient for coronal heating.}  { ",
        "introduction": "The idea that nanoflares are sufficiently numerous to maintain coronal temperatures was first proposed by Parker (1988). This theory implies that the coronal background emission is the result of nanoflares occurring continually throughout the solar atmosphere. Hence, it is probably unfeasible to observe these flares individually. This perhaps explains why observational studies have failed to agree on the importance of nanoflares with regard to coronal heating (Krucker \\& Benz 1998; Parnell \\& Jupp 2000; Aschwanden \\& Parnell 2002; Parnell 2004). The impracticality of nanoflare detection has motivated the development of models that attempt to show how coronal loops might dissipate their magnetic energy in the form of discrete heating events and thereby replenish coronal heating losses. The primary dissipation mechanism is thought to be magnetic reconnection, since strong evidence of this process has been found in observations of large-scale flares (Fletcher 2009; Qiu 2009). Coronal heating requires fast dissipation of magnetic energy, a requirement that is compatible with reconnection timescales.  A loop's excess magnetic energy is introduced via the random convective motions that occur at or below the loop's photospheric boundaries (i.e., footpoints). There are two possibilities for how the kinetic energy of these motions is dissipated within the loop (Klimchuk 2006). Direct current (DC) heating occurs when the Alfv\\'en time is small compared to the timescale of photospheric turbulence. Thus, the loop moves quasi-statically through a series of force-free equilibria until the accumulating magnetic stresses are dissipated as heat. If the Alfv\\'en time is slow compared to photospheric motions, alternating current (AC) heating takes place: footpoint motions generate waves (acoustic and MHD), some of which may deposit energy within the coronal loop. The waves that can penetrate the corona are thought to be restricted by type and frequency (Narain \\& Ulmschieder 1996; Hollweg 1984). Whether or not these waves carry sufficient energy to heat the corona is not certain (Porter et al. 1994). There is considerably less doubt, however, over the sufficiency of DC heating, which is the basis for the model presented here. (AC heating may still be partly reponsible for coronal heating: the reconnection mechanism mentioned above may generate waves within the corona.)  The coronal part of the loop is defined by its magnetic field (plasma beta, $\\beta$\\,$\\approx$\\,0.01) and, since the loop's magnetic flux and plasma are frozen together, its magnetic field is continually twisted by random swirling motions at the photosphere. If the driving is slow compared to Alfv\\'en timescales, the loop's force-free state can be represented by $\\nabla\\times\\vec{B}\\,=\\,\\alpha (\\vec{r})\\vec{B}$, where $\\alpha\\,=\\,(\\mu_{0}\\vec{j}\\cdot\\vec{B})$/$(|\\vec{B}|^{2})$ is the ratio of current density to magnetic field and \\textit{$\\vec{r}$} is a position vector (Woltjer 1958).  Many 3D MHD models have shown how such coronal loops exhibit current sheet formation during the nonlinear phase of an ideal kink instability (Baty \\& Heyvaerts 1996; Velli et al. 1997; Arber et al. 1999; Baty 2000). Essentially, helical current sheets become the site of Ohmic dissipation, resulting in a heating event. Further simulations have revealed the appropriate correlation between magnetic energy and Ohmic heating (Browning \\& Van der Linden 2003; Browning et al. 2008; Hood et al. 2009). The energy released by an instability depends on the complex dynamics of the magnetic reconnection events that occur inside the loop. Obviously, this is difficult to model - although this has been achieved by 3D MHD simulations. The computational expense of these simulations rules them out as a means of exploring fully the relationship between the $\\alpha$-profile and the amount of energy released. Fortunately, there is another way to calculate the energy release.  Relaxation theory states that when a magnetic field reaches instability (or is otherwise disrupted) it will evolve towards a minimum energy state such that the total magnetic axial flux and the \\textit{global} magnetic helicity are conserved (Taylor 1974, 1986). The relaxed state is the well known constant-$\\alpha$ or linear force-free field: \\begin{eqnarray} \\label{forcefree_field} \\nabla\\times\\vec{B} & = & \\alpha\\vec{B}. \\end{eqnarray} The original intention of this theory was to explain laboratory plasma phenomena; but latterly, it has been frequently applied to the solar corona (Heyvaerts \\& Priest 1984; Browning et al. 1986; Vekstein et al. 1993; Zhang \\& Low 2003; Priest et al. 2005). The helicity measures the self-linkage of the magnetic field, see Berger (1999). A modified expression for this quantity needs to be used since the field lines cross the photospheric boundaries, which creates non-zero normal flux and therefore removes gauge invariance. \\begin{eqnarray} \\label{relative_helicity} K & = & \\int_V (\\vec{A}+\\vec{A'})\\cdot (\\vec{B}-\\vec{B\\,'})\\,\\,dV, \\end{eqnarray} where \\textit{$\\vec{A}$} is the magnetic potential, \\textit{$\\vec{B\\,'}$} is the potential field with the same boundary conditions and \\textit{$\\vec{A'}$} is the corresponding vector potential. This relative helicity (Berger \\& Field 1984; Finn \\& Antonsen 1985) is the difference between the helicity of the actual field and that of a potential field (represented by the dash terms) with the same normal flux distribution. In ideal MHD, the helicity of every flux region is conserved. Taylor proposed that in the presence of reconnection all such local invariants are destroyed, whilst the global helicity is conserved. Helicity is still affected by global resistive diffusion, however, it can still be treated as a conserved quantity, if the change in helicity is substantially smaller than the change in magnetic energy. The results of one aforementioned MHD simulation (Browning et al. 2008) have shown that \\textit{$\\delta$K/K}\\,$\\sim$\\,10$^{-4}$ and \\textit{$\\delta$W/W}\\,$\\sim$\\,10$^{-2}$ (it should be noted that these figures are limited by the coarseness of the numerical grid used in the simulation). During relaxation, helicity is redistributed more evenly within the loop; as a consequence $\\alpha$ becomes invariant with radius. These properties enable one to locate a loop's relaxed state and hence calculate the energy released during relaxation, i.e., the upper limit of the flare energy.  In applying relaxation theory to coronal heating, it became clear that the effectiveness of the heating depends on how much free energy is stored before a relaxation event occurs (Heyvaerts \\& Priest 1984). Browning \\& Van der Linden (2003) proposed that relaxation is triggered by the onset of ideal MHD instability. Thus, the coronal field evolves quasi-statically in response to slow photospheric driving, building up free magnetic energy - until the field becomes unstable (i.e., the threshold for linear instability is reached). At this point, a dynamic heating event ensues. It should be noted that ideal instabilities are relevant (as opposed to resistive instabilities) because the time-scales are sufficiently fast. Dissipation and energy release can occur during the nonlinear phase of the instability, as has been demonstrated extensively by numerical simulations of the nonlinear kink instability (Galsgaard \\& Nordlund 1997; Velli et al. 1997; Lionello et al. 1998; Baty 2000; Gerrard et al. 2001). The helical deformation of the kink instability generates current sheets in the nonlinear regime, in which fast magnetic reconnection rapidly dissipates magnetic energy. Recently, it has been demonstrated using 3D MHD simulations that this process causes the field to relax towards a state which is closely-approximated as a constant-$\\alpha$ state, and the energy release is in good agreement with relaxation theory (Browning et al. 2008; Hood et al. 2009). These papers have investigated field profiles, based on the model of Browning \\& Van der Linden (2003), in the unstable region of parameter space and have shown that fast reconnection develops in a current sheet, with dissipation of magnetic energy and relaxation to a new, lower-energy state. Efficient \"mixing\" of the $\\alpha$ profile is facilitated in the later phases of the evolution, during which the current sheet fragments (Hood et al. 2009).  Supported by this evidence, the model developed here is based on the idea that coronal magnetic fields respond to photospheric driving by evolving through force-free equilibria, with a heating event triggered whenever the ideal instability threshold is reached. Following Browning \\& Van der Linden (2003), we use a simple cylindrical field model, and represent the nonlinear force-free field using a two-parameter family of current profiles in which $\\alpha$ is a piecewise constant. Browning and Van der Linden considered only individual heating events. The primary aim of this paper is to predict a \\textit{distribution} of heating events, generated by an ongoing stress-relax cycle. The photospheric driving stresses the coronal magnetic field until it becomes unstable and relaxes; then the driving resumes, and the process repeats, leading to a series of heating events as expected in the nanoflare coronal heating scenario. This is modelled through a monte-carlo approach, with the photospheric footpoint motions treated as random. Recent observational evidence (Abramenko et al. 2006) lends support to the idea that  random, turbulent photospheric motions can provide an energy source for coronal heating in Active Regions.  In calculating the relaxation and distribution of heating events, we need to know the linear instability threshold for line-tied coronal loops over a family of field profiles. As a consequence, we obtain some interesting new results concerning stability properties of cylindrical loops, for a much more extensive range of current profiles than previously studied (including profiles for which the sense of the twist of the field varies across the loop). A secondary aim is thus to explore the properties of the linear kink instability on this family of fields, and in particular, to determine the extent to which stability can be defined by any single quantity such as \"critical twist\".  The paper is structured in the following manner. Section 2 describes the composition of the loop model, along with the equations used to express the loop's magnetic field. The calculation of the loop's instability threshold is also explained, as well as the procedure for allowing the loop to repeatedly undergo instability and the equations used to determine the energy release associated with each relaxation. Section 3 outlines the analysis of possible critical parameter values for the onset of loop instability. The results of the simulations of energy release are presented in Sect. 4 as nanoflare energy distributions. Finally, in the last section, the results are discussed and our conclusions are given. ",
        "conclusions": "We have used a model based on relaxation theory to determine the energy release of a sequence of heating events generated by random photospheric footpoint motions. This is a \"stress-relax\" scenario, in which free energy is continually built-up in the corona due to photospheric driving - repeated relaxations occur whenever a critical condition is reached. We postulate that heating events are triggered by the onset of an ideal kink instability. Energy release then occurs by fast magnetic reconnection in current sheets generated during the nonlinear phase of the instability, and this is calculated by assuming a helicity-conserving relaxation to a minimum energy state. The energy release of any individual heating event depends on the current profile at the point where the instability threshold is crossed - thus heating events of a range of sizes are found.  The initial stressed fields are modelled using piecewise-constant profiles of $\\alpha$, yielding a two-parameter family of current profiles. Photospheric driving thus yields a random walk of the current profile. In order to pursue the calculation, the ideal instability threshold of a line-tied loop has been determined. The modelled current profiles incorporate regions of positive and negative twist as well as monotonic twist profiles, as have been more commonly studied. We have shown that it is not possible to determine the onset of instability by any simple criterion such as achieving a critical twist or critical helicity. Nevertheless, Fig. \\ref{rdtwpf} suggests that the location of the instability is coincident with the peak twist (the largest absolute twist).  A distribution of heating events has been calculated, both for the case of a single loop which undergoes repeated stressing and relaxation, and for an ensemble of identical loops which are randomly stressed (and for intermediate cases). For a sufficiently large number of events, a statistically-stable distribution of event sizes is obtained. As expected, the smallest events are the most common, and there is (as noted by Browning \\& Van der Linden 2003) a minimum event size for given loop parameters, perhaps corresponding to an \"elemental nanoflare\". More surprisingly, there is a second peak of event frequency at intermediate magnitudes, although this is somewhat reduced in size if an ensemble of loops is considered. This can be explained as follows. The first peak (at minimum energy) occurs simply because the range of energies near the minimum naturally encompasses the largest part of the instability threshold curve (because of the flatness of the minimum). The second peak is found because the instability threshold is most likely to be crossed in the part near to the region of constant-$\\alpha$.  An ensemble of identical loops - in which each individual loop starts from a randomly chosen initial state and is stressed until it undergoes a single relaxation event - is investigated in more detail. The distribution of heating events is qualitatively similar, although, the secondary peak of event frequency is much lower, and the decay of frequency with increasing energy is much flatter. A significant result is that, although the distribution is not a simple power law, the high-energy part of the distribution \\textit{is} well approximated by a power law with an index of around -8.3, which is considerably steeper than the minimum required for nanoflare heating to be effective (-2).  The model requires that the field is sometimes unstable - although most of the time, the field profile will be well within the stable region. Typically, the dimensional $\\alpha$ value is of magnitude 1\\,-\\,2, leading to dimensional values of $\\alpha$\\,$\\approx$\\,1\\,-\\,2\\,Mm$^{-1}$ (for a loop radius of 1 Mm). Note that this is the maximum value which could be found, and usually we would expect lower values. This is consistent with observations; for example R\\'egnier \\& Priest (2007) find $\\alpha$ magnitudes around 1\\,Mm$^{-1}$. Furthermore, a consequence of the fact that the fields are predicted to fluctuate between stable and unstable states is that we predict a value for the average horizontal field component (on average, this will be somewhat less than the value at marginal stability). At threshold, we find \\textit{B}$_\\perp$/\\textit{B}$_z$ is around 0.5, which means that its average value should be around 0.25. This agrees very well with the limit on this quantity required by Parker (1988), in order for the Poynting flux from the photosphere to match coronal heating requirements. Hence, we have an explanation for the critical shear value predicted by Parker. This also implies that the average heating flux from our model will be sufficient for coronal heating. Indeed, the fluxes derived from our results agree (especially if we raise the correlation time for photospheric motions) with the required values.  Relaxation theory makes no prediction about the spatial distribution of energy dissipation. However, this can be determined from numerical simulations. Recently, Hood et al. (2009) have shown that heating is well-distributed across the loop volume, as the current sheet, associated with the nonlinear kink instability, stretches and fragments, thereby filling the loop cross section. This has implications for the observed emission.  A further important consideration is that we study a single loop of fixed dimensions, with repeated stressing and relaxation. In practice, coronal heating events will occur in regions of field of varying sizes. A large flare will inevitably involve a large magnetic field volume, whereas nanoflares may involve small sets of fieldlines. This could be accounted for in our model by allowing the loop radius to be randomly distributed also; the effect would be to convolve a set of our energy distributions. A second limitation of the calculation is that we have assumed a uniform random walk in $\\alpha$ space. This is not realistic, since photospheric twisting motions are likely to be correlated across the loop cross section and therefore changes in which the twist in the outer later is similar to that in the core are much more likely: in other words, there should be a positive correlation between changes in $\\alpha_1$ (the core current) and changes in $\\alpha_2$ (the current in the outer layer), rather than these being independent random variables as assumed here. This will be considered further in future work.  As well as the improvements mentioned above, the model will be enhanced such that there is, more realistically, zero net-current carried by the loop. At present, in most cases the potential layer outside the loop contains azimuthal field. A current neutralisation layer will be inserted between the loop and the envelope - external magnetic fields will have an axial direction only (Hood et al. 2009), representing a response to localised photospheric twist motions.  The model has included a number of simplifications, and at this stage is more of a \"proof of principle\" showing that a distribution of heating events can be produced from an (almost) \\textit{ab initio} coronal heating model. These initial results demonstrate the viability of the model for further research.\\\\\\\\  \\small{\\textit{Acknowledgements}. We are grateful to Jim Klimchuk for helpful comments and M. R. B. thanks STFC for studentship support.}  \\appendix"
    },
    "1005/1005.1230_arXiv.txt": {
        "abstract": "We present the results of a spectroscopic investigation of two novel variable bright blue stars in the SMC, OGLE004336.91-732637.7 (\\va) and the periodically occulted star \\vbb~ (\\vb), whose photometric properties were reported by Mennickent et al. (2010). High-resolution spectra in the optical and far-UV show that both objects are actually A + B type binaries. Three spectra of \\vb~ show radial velocity variations, consistent with the photometric period of 184.26 days found in Mennickent et al. 2010.  The optical spectra of the metallic lines in both systems show combined absorption and emission components that imply that they are formed in a flattened envelope.  A comparison of the radial velocity variations in \\vb~ and the separation of the $V$ and $R$ emission components in the H$_{\\alpha}$ emission profile indicate that this envelope, and probably also the envelope around \\va, is a {\\it circumbinary} disk with a characteristic orbital radius some three times the radius of the binary system. The optical spectra of \\va~ and \\vb~ show, respectively, He\\,I emission lines and discrete Blue Absorption Components (``BACs\") in metallic lines. The high excitations of the He\\,I lines in the \\va~ spectrum and the complicated variations of Fe\\,II emission and absorption components with orbital phase in the spectrum of \\vb~ suggests that shocks occur between the winds and various static regions of the stars' co-rotating binary-disk complexes.  We suggest that BACs arise from wind shocks from the A star impacting the circumbinary disk and a stream of former wind-efflux from the B star accreting onto the A star. The latter picture is broadly similar to mass transfer occurring in the more evolved (but less massive)  Algol (B/A + K) systems, except that we envision transfer occurring in the other direction and not through the inner Lagrangian point.   Accordingly, we dub these objects prototype of a small group of Magellanic Cloud wind-interacting A + B binaries. ",
        "introduction": "Mennickent et al. (2002, hereafter M02) have reported the existence of a number of bright  blue stars in the Small Magellanic Cloud the light curves of which exhibit periodic or quasi-periodic variability in their OGLE (Udalski et al. 1997) $I$-band light curves.  In an initial follow up of the investigation, Mennickent et al. (2006; ``M06\") noted that some stars exhibit peculiar spectroscopic properties and multiple periods.  They gave initial estimates of spectral classifications based on low dispersion spectra, but they noted conflicting properties. In a companion study to this one, Mennickent et al. (2010; hereafter ``M10\") have selected two novel Type\\,3 variables\\footnote{ Mennickent Type-3 variables are SMC Be star candidates with $I$-band light curves varying periodically or quasi-periodically.} with novel properties named \\vaa~  ($\\equiv$ SMC-SC3-63371, MACHO ID 213.15560; hereafter \\va) and \\vbb~ ($\\equiv$ SMC-SC4-67145, MACHO ID 212.15735.6; hereafter \\vb). These Type\\,3 stars were chosen according to the sole additional arbitrary criterion that they are brighter than $m_v$ = 14.2. The photometric analysis in M10 was based on new OGLE III  light curves and found that \\va~ exhibits a double-minimum period with a period of 238.1 days and a second quasi-period that fluctuates between two values each near 15.35 days. The appearance of the longer period, which exhibits minima at unequal spacings and unequal depths suggests the appearance of an ellipsoidal variable. The high orbital eccentricity implied by these properties is confirmed by the inference from the sharp and unequal light maxima that reflection effects from a secondary occur during brief periastron passages. The light curve of \\vb~ reveals surprising periodic deep eclipses every 184.26 days, modulated over a superperiod of several cycles, and chaotic variability indicative of multiple shorter periods.   M06 and M10 have open a variety of questions  that can be pursued with high resolution spectra, which we undertake to pursue now. We have already remarked as one example that  the spectra of \\va~ and \\vb~ spectra are unusual in that they show properties of rather conflicting spectral types (strong Na\\,D  and strong H$_{\\alpha}$ and H$_{\\beta}$ emission components). Therefore, we obtained three high-dispersion optical spectra of \\va~ and \\vb~ during the period 2002-9. To obtain a better understanding  of how the spectra appear at short wavelengths we obtained a {\\it Far Ultraviolet Spectroscopic Explorer} ({\\it FUSE}) spectrum of both stars. These spectra and the {\\it FUSE} data form the basis of this investigation. The goal of our study is to elucidate the properties of the circumstellar environment of these objects and to pave the way to understanding the evolutionary state of this potential new subclass of Ae stars.       \\begin{table*} \\centering \\caption{Summary of observations. The MJD numbers at mid exposure are given. The periods detected in photometric data are also given. Phases refer to the ephemeris given in Mennickent et al. (2010).  CD refers to cross-disperser.} \\begin{tabular}{@{}cccccccccccc@{}} object &periods (d) & instrument &UT-date   &    airmass &$\\Delta\\lambda$ (\\AA)&   grating & exptime (s) &mjd-obs &phase &$S/N$\\\\ \\hline \\va~ & 238, 15 &UVES &2002/05/17&  1.96&3100--8500 &CD\\#1,3&  600 &52411.39633  &0.28&25\\\\ \\va~ & &UVES &2002/05/17&  1.91&3100--10400  &CD\\#1,4&  600 & 52411.40536 &0.28&20\\\\ \\vb~ & 184 &UVES &2002/05/17&  1.86&3100--8500  &CD\\#1,3&  600 &52411.41614 &0.23&20\\\\ \\vb~ & &UVES &2002/05/17&  1.81&3100-10400  &CD\\#1,4&  600 & 52411.42782 &0.23&15\\\\ \\va~ & &MIKE  &2007/11/09&  1.43 &3390--9410 &echelle&  500 &54413.03654  &0.69   & 25        \\\\ \\vb~ & &MIKE  &2007/11/09&  1.42&3390--9410 &echelle&  500 &54413.04421   &0.10    &  40      \\\\ \\va~ & &Echelle  &2009/08/25  &  1.44 &3940--7490 &echelle&  2000 &55068.24680 &0.44   &      17 \\\\ \\vb~ & &Echelle &2009/08/25&  1.40&3940--7490 &echelle&2000   &55068.27751    &0.66    &17       \\\\ \\hline \\end{tabular} \\end{table*}    \\begin{table*} \\centering \\caption{Summary of {\\it FUSE} Observations.  Phases refer to the ephemeris given in Mennickent et al. (2010).} \\begin{tabular}{@{}cccccc@{}} object & UT-date  &UT-start & exptime (s) &mjd-obs &phase\\\\ \\hline \\va~   &  2006/10/05&04:05:06& 2430  &54013.18427083 &0.01\\\\ \\vb~   &  2006/11/30&11:36:26& 2736  &54069.49946759 &0.24\\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "Even as luminous variables in the SMC, our two program stars exhibit a extraordinary set of photometric and spectroscopic properties. These include a multiperiodic photometric variability, a $\\approx$B3 type spectrum in the FUV, an optical A-type spectrum with Balmer emission lines and metallic lines from an intervening disk. For SMC-SC4 the spectral anomalies include the presence of one and sometimes two Discrete Blue Absorption Components (BACs) in metallic lines and RV variations. We have remarked that the optical spectrum of the SMC supergiant N82 in the SMC exhibits BACs and strong H$_{\\alpha}$ emission. However, unlike N82, neither of our stars shows forbidden lines. Because UV spectra are not available for N82, nor an optical RV campaign mounted, its suggested binary nature has not been tested.  We argue that our program stars reflect a comparatively brief stage in the life of intermediate mass $\\sim$ 9 M$_{\\sun}$ binary components. With separations of a few AU between the components, the winds of the two stars can interact and produce a large array of spectroscopic phenomenology. This phenomenology can include emissions in lines of intermediate and/or highly excited ions (e.g., He\\,I) and time-variable pattern of sharp absorptions, in addition to the BACs. It is an open question why as many as two objects have such similarly unusual properties from a sample of only bright eight Mennickent Type\\,3 variables.  In any case we may dub these objects prototypes of a presumably small group of Magellanic Cloud wind-interacting A + B binaries.    We speculate that the emission components originate from the interaction of that component of the A and B star winds in the zone near the line that intersects the centers of the two stars. The double peaked emissions in the Balmer lines disclose an emitting  Keplerian disc, which from the peak separation has a larger extent than the binary separation. We have suggested, in part because of the He\\,I line emissions in the \\va~ spectrum, that wind-wind interactions cause localized heating and ionizing photons capable of exciting the observed level of H$_{\\alpha}$ model.  The Fe\\,II line emission strengths are {\\it anticorrelated} with the line's oscillator strengths. We have attributed this extraordinary circumstance to their formation in a geometrically thin but optically thick and two layered column, where emission is formed in the deeper hotter shock and the absorption is just outside this region (though in the foreground of the observer's sight line). We have also speculated that wind-wind interactions in the zone between the stars collisionally excite He\\,I line emission and also produce Balmer emission in the quasi-Keplerian circumbinary disc.  We have surmised that the primary BAC (absorption) feature is caused by the violent impact of the expanding A star wind into this CB disk. However, although we do not have enough observations to characterize the time-dependent secondary components well, it appears that the strengths of the emission components in either one of the far wings of the 3\\,eV Fe\\,II lines are physically associated with absorptions in the opposite far wing.  We have speculated that a rich line profile variability, including the so-called secondary BACs and emission components, is a consequence of a matter stream originating from the wind of the B star and settling on the A star.  This is a reverse Algol scenario in which matter is not  constrained to flow only through the inner Lagrangian point. Perhaps this settling processes is responsible for the eclipses in the continuum light curve of SMC-SC4.    A number of observing programs can be suggested to test these ideas and conceivably offer new alternatives:  \\begin{itemize}  \\item A high-resolution monitoring of these objects should be undertaken to construct a radial velocity curve and determine the periods, separations, and inclinations of these systems. These same data can be used to trace the velocities of the various components in the primaries' spectra. For \\vb~ the features of interest are the BACs and the pairs of absorptions and emissions in the wings of the Fe\\,II lines. For \\va~ it is important to determine whether the velocity of the He\\,I lines is constant or shows a small amplitude variation around the system's velocity.  \\item UV observations of the prominent resonance lines are needed to characterize the wind strength, velocity, and its behavior through the orbital cycle. We expect that these features are diagnostics of the B star's wind. It may also be possible to trace velocity changes in signatures of the A star's wind, such as an arguably Ca\\,II K emission feature (Fig. 4b).  In general, we expect the profiles of the UV resonance lines to be highly complex and variable.  \\item Any understanding of these particular objects rests on the expansion and adequate definition of this class, which in term entails finding more members. So far our small sample is drawn from a group of luminous optical/IR variables in the SMC. One prong of this search would be to find stars whose spectra exhibit BACs. To date we have obtained moderate resolution spectra of a number of luminous B stars in the OGLE surveys of both Clouds. We find that two LMC stars have properties similar to our two prototype objects. These are OGLE05141821-6912350 and OGLE00552027-7237101. In addition, Mennickent et al. (2010) noted that the object OGLE LMC LPV-41682 (P = 219.9 days; Soszynski 2010) exhibits modulated eclipses similar to those of \\vb.  Spectra of the two first objects display a range of H$_{\\alpha}$ emission profile intensities, strong absorptions among high level lines, and variable (possibly regular) light curves. Although the resolution of these spectra is sufficient to show at least strong forbidden emission lines, none are visible. If many of these stars are members of the group exemplified by \\va~ and \\vb, it will indicate that their properties are not isolated to a particular evolutionary age and/or angular momentum state. This is a possibility that we cannot yet dismiss.  \\end{itemize}"
    },
    "1005/1005.4415_arXiv.txt": {
        "abstract": "As the quality and quantity of astrophysical data continue to improve, the precision with which certain astrophysical events can be timed becomes limited not by the data themselves, but by the manner, standard, and uniformity with which time itself is referenced. While some areas of astronomy (most notably pulsar studies) have required absolute time stamps with precisions of considerably better than 1 minute for many decades, recently new areas have crossed into this regime. In particular, in the exoplanet community, we have found that the (typically unspecified) time standards adopted by various groups can differ by as much as a minute. Left uncorrected, this ambiguity may be mistaken for transit timing variations and bias eccentricity measurements. We argue that, since the commonly-used Julian Date, as well as its heliocentric and barycentric counterparts, can be specified in several time standards, it is imperative that their time standards always be reported when accuracies of 1 minute are required. We summarize the rationale behind our recommendation to quote the site arrival time, in addition to using \\bjdtdb, the Barycentric Julian Date in the Barycentric Dynamical Time standard for any astrophysical event. The \\bjdtdb is the most practical absolute time stamp for extra-terrestrial phenomena, and is ultimately limited by the properties of the target system. We compile a general summary of factors that must be considered in order to achieve timing precisions ranging from 15 minutes to 1 $\\mu$s. Finally, we provide software tools that, in principal, allow one to calculate \\bjdtdb \\ to a precision of 1 $\\mu$s for any target from anywhere on Earth or from any spacecraft. ",
        "introduction": "Precise timing of astrophysical events is one of the fundamental tools of astronomy, and is an important component of essentially every area of study. There are two basic sources of uncertainty in timing: the astrophysical data characterizing the event, and the time stamp with which the event is referenced. Unfortunately, since the accuracy of the time stamp is something that is often taken for granted, the improvements in data are sometimes not accompanied (or are not uniformly accompanied) by the requisite improvements in accuracy of the time stamp used. This situation can lead to confusion, or even spurious inferences.  Timing plays a particularly important role in the study of exoplanets. Indeed, many of the ways in which exoplanets are discovered involve the detection of transient or time-variable phenomena, including the radial velocity, transit, microlensing, and astrometry techniques. Furthermore, in some cases much can be learned about planetary systems from the precise timing of these phenomena. As examples, the measurement of terrestrial parallax in microlensing allows one to infer the mass of the primary lens and so the planetary companion \\citep[e.g.,][]{gould2009}, and one can constrain the eccentricity of transiting planets by comparing the times of primary transits and secondary eclipses \\citep[e.g.,][]{deming05,charbonneau05,knutson07}. Possibly the most promising application of timing in exoplanets, however, comes from transit timing variations (TTVs). With an exquisitely periodic phenomena like transiting planets, we will be able to measure many effects using the departures from strict periodicity, such as the gravitational perturbations from additional planets \\citep{miralda02,holman05,agol05}, trojans \\citep{ford06b,ford07}, and moons \\citep{kipping09}, stellar quadrupoles \\citep{miralda02}, tidal deformations, general relativistic precession \\citep{jordan08,pal08a,heyl07}, orbital decay \\citep[e.g.,][]{sasselov03}, and proper motion \\citep{rafikov09}.  Because of their great potential, TTVs have become the focus of many groups. The typical data-limited transit timing precisions of most observations are around 1 minute, with the best transit time precision yet achieved of a few seconds \\citep{pont07}. However, as discussed above, accuracies of transit times are limited not only by the data themselves, but also by the time stamp used. In order to make these difficult measurements useful, it is critical that a time stamp be used that is considerably more accurate than the uncertainty due to the data themselves. Furthermore, since thorough characterization of TTVs will require the use of all available data spanning many years from several groups, this time stamp must be stable in the long term, and all groups must clearly convey how it was calculated.  Unfortunately, we have discovered that, in the exoplanet community, the Julian Date (JD) and its geocentric (GJD), heliocentric (HJD), and barycentric (BJD) counterparts are currently being quoted in several different, and often unspecified, time standards. In addition, the site arrival time and its time standard is not quoted. This general lack of homogeneity and specificity leaves quoted time stamps ambiguous at the 1 minute level. More alarmingly, the most commonly-used time standard, the Coordinated Universal Time (UTC), is discontinuous and drifts with the addition of each leap second roughly each year.  The pulsar community has solved the problem of precise timing well beyond the level that is currently necessary for exoplanet studies, and we can benefit from the techniques they have developed over the past 40 years. In particular, their current state of the art program ({\\tt TEMPO2}) models pulsar arrival times to 1 ns precision \\citep{hobbs06,edwards06}. This program is highly specialized and generally cannot be applied outside of pulsar timing observations, but many of the effects they consider are relevant to optical observers in the exoplanet community.  In this article, we summarize the effects one must consider in order to achieve timing accuracy of 1 $\\mu$s -- well beyond the accuracy that will likely be required by the exoplanet community for the foreseeable future.  Section \\ref{theory} provides the background required to understand each of the effects that could change the arrival time of a photon. They are listed in order of decreasing magnitude, so latter subsections can be ignored for low-precision measurements.  Section \\ref{practice} discusses the practical limitations to achieving high-precision timing. We begin with the effects which may cause errors that are comparable to or exceed the BJD correction. These should be read and understood by everyone. We continue with remaining effects, in order of decreasing magnitude, which can be ignored for low-precision ($> 30$ ms) measurements. We conclude \\S \\ 3 by listing additional effects, the errors due to which are negligible ($< 1\\mu$s).  We begin \\S \\ \\ref{sec:calculating} by detailing the procedure one must follow in order to calculate the \\bjdtdb, which is designed to be a useful reference for those already familiar with the concepts of precision timing. In the latter part of this section, we describe our particular IDL and web-based implementation of this procedure.  Lastly, in the Appendix, we discuss some of our specific findings about the time stamps currently in use and how these are calculated throughout the exoplanet community.  While we focus on the effects of timing on the optical/infrared exoplanet community, timing precision of order 1 minute is necessary for many other areas, such as the study of rapidly rotating white dwarfs \\citep{euchner06}. This article should be equally applicable in such cases. ",
        "conclusions": "Timing of transient events is a powerful tool for characterizing many astronomical phenomena. In the field of exoplanets in particular, the search for variations in the times of primary transits and secondary eclipses, or transit timing variations, is one of the most promising new techniques for studying planetary systems.  The accuracy with which transit times, and indeed any transient phenomenon, can be measured is limited not only by the data themselves, but by the time stamp to which the transit time is referenced. As the quality of transit timing data crosses the threshold of 1 minute precision, the precise time standard and reference frame in which event times are quoted becomes important. Achieving uniform and accurate time stamps with accuracies of better than 1 minute that can reliably be compared to one another requires extraordinary care in both our techniques and our terminology. We have found that the time standards adopted by various groups that measure transit times can differ by as much as a minute, and are typically left unspecified. As these ambiguities can be significant compared to the timing precisions that are quoted, they may therefore lead to spurious detections of transit timing variations or biased eccentricity measurements.  Here we have summarized the effects one must consider in order to achieve timing precision of 1 $\\mu$s. We argue that the \\bjdtdb \\ is nearly the ideal time stamp, being as reliable as any time stamp can be without being unique to each target system. On the other hand, \\bjdutc \\ and the HJD in any form should be avoided whenever possible. Most importantly, we emphasize that the time standard should always be explicitly stated. Any time stamp that is quoted without a time standard should be assumed to be uncertain to at least 1 minute. Unless the time standards used in programs or algorithms have been independently confirmed, one should avoid using ones that do not precisely specify the input and output time standard.  In addition, the arrival time at the observing site along with its time standard (e.g., \\jdutc) should also be specified. This will remove any ambiguity in the time stamp, allow others to apply improved corrections should more precise ephemerides become available in the future, and allow others to check that original conversion was done accurately enough for their purpose.  Finally, we have written an IDL program for general use that facilitates the use of \\bjdtdb \\ to an accuracy of 1 $\\mu$s, provided that the inputs are sufficiently precise, and we provide a web-based interface to its most useful features."
    },
    "1005/1005.1837_arXiv.txt": {
        "abstract": "VisIVO \u00a0is an integrated suite of tools and services specifically designed for the Virtual Observatory. This suite constitutes a software framework for effective visual discovery in currently available (and next-generation) very large-scale astrophysical datasets. VisIVO consists of VisiVO Desktop - a stand alone application for interactive visualization on standard PCs, VisIVO Server - a grid-enabled platform for high performance visualization and VisIVO Web - a custom designed web portal supporting services based on the VisIVO Server functionality. The main characteristic of VisIVO is support for high-performance, multidimensional visualization of very large-scale astrophysical datasets. Users can obtain meaningful visualizations rapidly while preserving full and intuitive control of the relevant visualization parameters. This paper focuses on  newly developed integrated tools in VisIVO Server allowing intuitive visual discovery with 3D views being created from data tables. VisIVO Server can be installed easily on any web server with a database repository. We discuss briefly aspects of our implementation of VisiVO Server on a computational grid and also outline the functionality of the services offered by VisIVO Web. Finally we conclude with a summary of our work and pointers to future developments. ",
        "introduction": "An essential part of astrophysical research is the necessity to employ computer graphics and scientific visualization tools for displaying appropriately multi-dimensional data plots and images either from real-world observations or from modern, highly-complex numerical simulations. The latest generation of graphics and visualization software tools offer the astrophysical community robust instruments for data analysis and advanced data exploration by using: \\begin{enumerate} \\item {\\it High Performance Computing} and {\\it Multithreading} to exploit multi-core CPUs and emerging powerful graphics boards, thus allowing real-time interaction with large-scale datasets.  \\item {\\it Interoperability} so as to allow different applications, each specialized for different purposes, to operate on shared datasets.  \\item {\\it Collaborative Workflows} so that several users can work simultaneously - perhaps at different geographical locations - on identical datasets exchanging information and visualization experiences.  \\end{enumerate}  Astrophysicists are currently witnessing an unprecedented growth in the quality and quantity of datasets coming from real-world observations and numerical simulations. For example, the increasing availability of high performance and grid computing facilities \\citep{BeccianiGrid} \\citep{ube} has given the possibility to perform numerical simulations of several dimensions. Also current, such as the Sloan Digital Sky Survey (SDSS) \\citep{SDSS} \\citep{WWWSDSS}, and next-generation sky surveys, such as the Low Frequency Array (LOFAR) \\citep{LOFAR}, the Square Kilometre Array (SKA) \\citep{SKA}, the Large Synoptic Survey Telescope System (LSST)\\citep{LSST}, and the Dark Energy Survey (DES) \\citep{DES}, are planned to collect large amounts of raw data eventually resulting in several hundreds of terabytes of data in tabular form. These developments place the emphasis on the problems associated with data post-processing. For storage, very large-scale distributed databases will be required and scientists will need to interact with them effectively using very high performance computer graphics and scientific visualisation tools as simply downloading will be prohibitively costly. The emerging need is not only to provide data analysis and exploration tools that can run on standard PCs satisfactorily, but also to offer appropriate interoperable tools that can run on servers and standard PCs simultaneously.\\\\ This paper describes VisIVO \u00a0Server \\citep{visivoserver}  a new platform for astrophysical visualization of large-scale datasets that can be easily installed on any computing server and can handle datasets in cooperation with VisIVO Desktop \\citep{wwwvisivodk}, our previously developed, stand-alone visualization application for standard PCs. Section 2 includes a short discussion on existing visualization tools for astrophysics and outlines the VisIVO Desktop functionality. The VisIVO Server is discussed in detail in section 3, including a presentation of its main components namely, VisIVO Importer, VisIVO Filters and VisIVO Viewer. Section 4 discusses aspects of our current grid installation of VisIVO Server while VisIVO Web, a portal for providing VisIVO Server services, is outlined in section 5. Finally section 7 presents a summary of our work and includes pointers to future developments. ",
        "conclusions": "VisIVO is an integrated suite of tools and services specifically designed for the Virtual Observatory. This paper focused on VisIVO Server, a new platform for astrophysical visualization of large-scale datasets that can be easily installed on any computing server. VisIVO Server is an open-source collection of data processing and visualisation modules allowing fast rendering of 3D views. VisIVO Server is founded on the functionality of VisIVO Desktop, our previously developed stand-alone astrophysical visualization application for standard PCs.The defining characteristic of VisIVO Server is that no restrictions are imposed on dimensionality of datasets.\\\\ We described the constituent parts of VisIVO Server in detail, namely VisIVO Importer, VisIVO Filter and VisIVO Viewer. VisiVO Importer converts user prescribed datasets into a highly efficient internal data format employed by VisIVO Filters and VisIVO Viewer. VisIVO Filters are a collection of several processing modules for constructing customised data tables from VisIVO Importer. Finally, VisIVO Viewer creates 3D views of astrophysical datasets, currently supporting points, volumes iso-surfaces and ray-tracing using a customized algorithm called Splotch.\\\\ We discussed different visualization scenarios and demonstrated example renderings. Users can obtain meaningful visualizations rapidly while preserving full and intuitive control of the relevant visualization parameters. We also described briefly a computational grid realisation of VisIVO Server and also discussed VisIVO Web - a custom designed web portal supporting services based on the VisIVO Server functionality. The VisIVO Server software is distributed under a GPL license  for non-commercial use only. It is an open source project and can be simply downloaded from the sourceforge code repository.\\\\ VisIVO Server and VisIVO Desktop can exchange data using the same internal data format (VBT). They can be closely integrated, but are also complementary and independent of each other. For example VisIVO Server can create a preview of a dataset or an appropriately-defined subsampled version, then VIsIVO Desktop can be employed for visualization and subsequent interaction.\\\\ Currently we are working towards an integrated visualisation functionality so that VisIVO Server and VisIVO Desktop are wrapped around a unified visualisation kernel. Our vision is to maintain VisIVO Server for high performance computing (e.g. grid environments, large memories) and VisIVO Desktop to support interactive front-end interfaces providing an advanced GUI and exploiting the full graphical capabilities of modern PCs (e.g. fast multi-processor architectures, emerging GPUs). Towards this aim we will need to define a full communication protocol between the Server and Desktop versions, so that applications for each will be possible to be developed independently."
    },
    "1005/1005.3765.txt": {
        "abstract": "Evidence for an accelerated expansion of the universe as it has been revealed ten years ago by the Hubble diagram of distant type Ia supernovae represents one of the major modern revolutions for fundamental physics and cosmology. It is yet unclear whether the explanation of the fact that gravity becomes repulsive on large scales should be found within general relativity  or within a new theory of gravitation. However, existing evidences for this acceleration all come from astrophysical observations. Before accepting a drastic revision of fundamental physics, it is interesting to critically examine the present situation of the astrophysical observations and the possible limitation in their interpretation. In this review, the main various observational probes are presented as well as the framework to interpret them with special attention to the complex astrophysics  and theoretical hypotheses that may limit actual evidences for the acceleration of the expansion.  Even when scrutinized with sceptical eyes, the evidence for an accelerating universe is robust. Investigation of its very origin appears as the most fascinating challenge of modern physics. ",
        "introduction": "\\label{intro} Modern cosmology has achieved remarkable progresses during the last fifty years. The  general picture originally designed as the \"Primeval atom\" by Lema\\^\\i tre and which has become the \"Big Bang\" model according to the word of one of its most famous opponent, F.~Hoyle, is now recognized as the successful scientific representation of the world at the large scales (in space and in time) we can measure\\footnote{It is fair to say that few scientists are still opposed to the \"Big Bang\" picture. Most of the serious opponents try to demonstrate that some observational facts, most often only one, which are coming  in support of the Big Bang may be interpreted in a different way and therefore the whole construction has to  be questioned.  It is useful to remember that the success of a scientific model is -- in some sense -- measured by the number of predictions it leads to and how many are successful. Newton theory of gravity is  wrong, but nevertheless it remains a high quality scientific theory because of its past (and present) successes. It is in this sense that modern cosmology should be regarded as successful, and this will remain in the future, even if it might be regarded as being \"wrong\"...}. Construction of this picture has necessitated the successive abandonment of philosophical and scientifi ideas, some of which are not only those of physicists but are also shared by the more general public. Maybe the first to be given up was the idea that it is hopeless to try to figure out a global picture of the universe.  Although heroic pioneers can be traced back since a long time ago, undoubtedly Einstein is the first one to directly address the question of handling the universe as a single physical object with an appropriate tool in hands, general relativity (GR).  The fact that one of the very first applications of his theory was  Cosmology is an evidence that the cosmological question was central in his thoughts. This is reinforced by the fact that he proposed to modify the initial formulation of his theory given the problems he encountered. It is commonly said that Einstein introduced the cosmological constant in order to obtain a static universe because he was reluctant to the idea of an expanding picture. This formulation is  inappropriate: it merely suggests that Einstein was willing to avoid an expanding universe, while he actually wanted to find at least one solution to the cosmological problem with an initial formulation which assumes the universe to be stationary, there is no indication in his seminal paper that he wished to reject a non-stationnary universe:   he actually started his paper by  discussing   the problem of having  mass at large distances in the Newtonian approach, noticing that this leads to divergence of the  potential. He insisted that this  would lead to unacceptably large velocities for stars. He also quoted that this can be cured by assuming a correcting term to the Newtonian potential  equation: \\begin{equation} \\nabla^2 {\\varphi} - \\lambda {\\varphi} = 4\\pi G \\rho \\label{} \\end{equation} and then proposed to modify his initial theory with the addition of the cosmological term $\\Lambda$. This allowed him to construct the first relativistic cosmological model, the Einstein solution, which is spatially closed (because being spherical) and static. In 1919, de Sitter discovered a new solution to Einstein equation which was written in a stationary form\\footnote{The choice of the coordinates system lead  to a form of the metric for which the coefficients are constant.} and contains no matter (but a non-zero cosmological constant). It is only a few years after, in 1925, that Lema\\^\\i tre identified the de Sitter solution to  an homogeneous  expanding universe \\citep{Lemaitre1925}. Friedmann \\cite{Friedmann1922,Friedmann1924} found the  general homogeneous solutions, providing the equation for the scale factor $R(t)$ and recognised their  expanding nature. It is somewhat surprising that his work has remained totally unnoticed, despite a controversy with Einstein. During this period  it is clear that the nature of the redshift of what Hubble had identified as extragalactic nebulae became a question addressed by many astronomers. Slipher's discovery was probably much more intriguing now that the nature of nebulae had been identified. Eddington is often mentioned as  the first astronomer to have noticed a possible connection between redshift and the de Sitter solution \\cite{Lemaitre1925}.  Carl Wirtz was clearly looking for a relation between redshift and distance \\cite{Wirtz1} and had in mind the possible cosmological information it might provide \\cite{Wirtz2}. Lema\\^\\i tre reestablished the equations Friedmann derived, showing that expanding solutions were leading to a redshift proportional to the distance. He proposed that this effect  was the origin of known redshift and provided the first estimation of the Hubble constant \\cite{Lemaitre1927}. In 1929,  Robertson published the now so called Robertson-Walker metric. In the same paper he mentioned that  distant sources appear to have a frequency shift.  %Historical view on cosmologcial constant dominated universe. ",
        "conclusions": "The Copernician model of the world  was the first revolution of a series in the construction of modern cosmology, and the discovery of the accelerated expansion being the latest in date. Theoretical considerations have always been a source of remarkable observational investigation and Cosmology has always benefited from the confrontation of models with observations. Since the thirties, the big bang picture, the modern version of Lema\\^\\i tre's primeval atom has been remarkably successful, based on simple assumptions and physics laws that have been validated  by accurate experimental results. Although alternative theories have been developed, these alternative were based on hypothetical unknown physics advocated to interpret cosmological observations. None of these alternative theories has produced predictions that have been comforted a posteriori. Rather new observations in agreement with predictions of the big bang picture necessitated deep revision of the unorthodox view, at the cost of rather ad hoc assumptions added to fit the new observations. The situation has evolved when the standard picture has necessitated the introduction of new ingredients, first dark matter and more recently dark energy. The very nature of these new ingredients, which are supposed to dominate the mean density of the universe has not been established by direct laboratory experiments, nor by astronomical observations, and this situation may some time lead to the question whether cosmologists have not  introduced new aethers. We had the opportunity to see that the situation is not so. The introduction of -cold- non-baryonic dark matter has led to specific predictions, the amplitude of the fluctuations of the cosmological background on various angular scales, which were verified with high accuracy precision. The presence of dark energy has lead to a specific prediction, the shape of the power spectrum on large scales, which has been verified  a posteriori. Although, the inclusion of a cosmological constant was concomitant to general relativity, the actual origin of dark energy remains totally unknown and the presence of dark energy in the present day universe represents probably the most fundamental and unexpected new element in modern physics.  %very different asumptions, but leading to similar observational predictions. %More precisely, it is fair to say that alternatives views have been build to %challenge the dominant paradigm, but at the cost of rather ad hoc asumptions.  % For tables use %\\begin{table} % table caption is above the table %\\caption{Please write your table caption here} %\\label{tab:1}       % Give a unique label % For LaTeX tables use %\\begin{tabular}{lll} %\\hline\\noalign{\\smallskip} %first & second & third  \\\\ %\\noalign{\\smallskip}\\hline\\noalign{\\smallskip} %number & number & number \\\\ %number & number & number \\\\ %\\noalign{\\smallskip}\\hline %\\end{tabular} %\\end{table}   %"
    },
    "1005/1005.3017_arXiv.txt": {
        "abstract": "Most stars are born in clusters and the resulting gravitational interactions between cluster members may significantly affect the evolution of circumstellar discs and therefore the formation of planets and brown dwarfs. Recent findings % suggest that tidal perturbations of typical circumstellar discs due to close encounters may inhibit rather than trigger disc fragmentation and so would seem to rule out planet formation by external tidal stimuli.  However, the disc models in these calculations were restricted to disc radii of 40~AU and disc masses below 0.1~\\tmsun.  Here we show that even  modest encounters can trigger fragmentation around 100~AU in the sorts of massive ($\\sim0.5\\,\\msun$), extended ($\\ge100\\unit{AU}$) discs that are observed around young  stars. Tidal perturbation alone can do this, no disc-disc collision is required. We also show that very-low-mass binary systems can form through the interaction of objects in the disc.  In our computations, otherwise non-fragmenting massive discs, once perturbed, fragment into several objects between about 0.01 and 0.1~\\tmsun, i.e., over the whole brown dwarf mass range. Typically these orbit on highly eccentric orbits or are even ejected.  While probably not suitable for the formation of Jupiter- or Neptune-type planets, our scenario provides a possible formation mechanism for brown dwarfs and very massive planets which, interestingly, leads to a mass distribution consistent with the canonical substellar IMF. As a minor outcome, a possible explanation for the origin of misaligned extrasolar planetary systems is discussed. ",
        "introduction": "The origin of brown dwarfs (BDs, to which we also include very low-mass stars $<0.1$\\tmsun) as well as the formation of massive exoplanets is contentious. The most common idea is that BDs form like stars, and planets form by accretion onto a core formed from dust conglomeration within a circumstellar disc. However, there is increasing evidence that an alternative formation scenario for BDs and, possibly, some of the most massive exoplanets is required. BDs differ from low-mass stars in several ways. Firstly, there is the `brown dwarf desert' -- a lack of BD companions to stars, especially at separations below about 30~AU within which companion planets and stars are common.  In addition, the properties of BD-BD binaries are very different from star-star binaries, showing a significant lack of systems wider than $\\approx 20$~AU beyond which most stellar binaries are found. Furthermore, while typically every second star is a binary system only every fifth BD is a binary. Clearly we must ask ourselves why stars and BDs are so different if they form by the same mechanism. Indeed, in order to numerically set-up a realistic stellar and BD population, stars and BDs need to be according to different algorithmic pairing rules \\citet{Ketal03}. In other words, they follow different formation channels.  The fragmentation of extended circumstellar discs is one of the most promising alternatives to the standard scenario of star-like formation for BDs \\citep{GoWi07}. Research on circumstellar disc evolution has made great progress in recent years, but there are still many unanswered questions \\citep{PapTer06,Henning2008,Hillenbrand2008}. One issue that has gained surprisingly little attention in the past is the role of gravitational interactions between the disc and passing stars within the young host stellar cluster. Since stars are generally born in clusters \\citep{FoRi09}, rather than as isolated objects, star-disc and disc-disc interactions \\citep{NSoPE1,NSoPE2,pfalzner:2005,TKT05} probably play an important role for the dynamical evolution of such discs and thus for massive planet and BD formation through fragmentation of the discs.  For a typical open star cluster hosting about 1000 stars within a half-mass radius of 0.5~pc, encounters closer than 500~AU will happen approximately every 10~Myr \\citep{TKT05}. They are even more frequent in denser clusters in which are thought to form at least half of the stars in the Galaxy. Encounters in star clusters are therefore expected to play a significant role in the evolution of extended discs of several 100~AU radii.  The formation of planets through core accretion is typically expected to happen within the innermost 100~AU of a circumstellar disc \\citep{Hillenbrand2008}, although there are many unanswered questions especially concerning the very first stages of core accretion \\citep{Henning2008}. Fragmentation has been proposed as an alternative way to form the most massive gas giant planets by \\citet{Mayeretal:2002,Boss97,2004ApJ...610..456B,Boss06}. In recent years, however, there has been a growing consensus that disc fragmentation cannot produce gas giant planets in the inner regions of discs for two reasons: (1) heating from the central star stabilizes the innermost regions and thus inhibits fragmentation, and (2) these regions cannot cool fast enough to allow temporary gas clumps to collapse \\citep{Lodatoetal2007,StaWhi08,Caietal2010}. This picture changes greatly beyond about 100~AU where cooling becomes sufficiently efficient for gravitational collapse to occur. Consequently, disc fragmentation has been proposed as an alternative mechanism for the formation of BDs \\citep{WhiSta06,GoWi07}. Most of these investigations, however, assume that the disc evolves in isolation.  Star-disc and disc-disc collisions have been investigated in previous studies \\citep{NSoPE1,NSoPE2,NSoPE3} but without a realistic treatment of the radiative heat transfer in the disc. From SPH computations \\citet{FoRi09} deduce that close encounters inhibit fragmentation in typical protoplanetary discs (assuming an initial radius of 40~AU) rather than inducing it. In contrast to that, a recent study of disc-disc collisions \\citep{Shenetal10} has, however, shown that the direct compression of gas by the close encounters of stars hosting massive extended discs (disc mass $\\approx 0.5$~\\tmsun, disc radius $\\approx 1000$~AU), may trigger the formation of substellar companions. However such massive, extended discs are probably short-lived \\citep{EiCa06} (in contrast to typical protoplanetary discs which are about ten times smaller and less massive), so that pairwise interactions of two discs of this kind are expected to be rare. The purely gravitational interaction of a low-mass star with no (or only a small) protoplanetary disc with a star hosting a massive disc would be much more likely. Although even these encounters may probably account for only a fraction of all BDs, especially in loose associations like Taurus-Auriga, they provide an additional channel of forming BDs out of discs that would otherwise probably dissolve without ever fragmenting. In this paper we address this issue. In Section~\\ref{sec:model}, besides an estimate of the encounter probability, the model of the circumstellar disc is described. Section~\\ref{sec:method} depicts the computational methods and gives an overview of the model parameters studied in this work. The results are then presented in Section~\\ref{sec:results} and discussed in Section~\\ref{sec:discussion} ",
        "conclusions": "\\label{sec:conclusions} In a series of SPH calculations we have shown that massive ($\\sim0,5\\,\\msun$), extended ($\\ge100\\unit{AU}$) circumstellar discs can be stable when isolated, but fragment when perturbed by a moderately close and slightly inclined stellar encounter. Binaries formed in two cases via triple encounters of companions and via grazing encounters of accreting envelopes. This agrees with binary formation in self-fragmenting discs \\citep{StaWhi09a}. We further found that the mass distribution of companions formed in discs is in agreement with the canonical substellar IMF as a separate population \\citep{Kr01,KB03b,TK07,TK08}. We have, however, to note that direct tidally induced fragmentation is probably only responsible for a fraction of all disc fragmentation events due to the relatively small orbital parameter subset that is suitable for fragmentation of discs of about 0.5~\\tmsun. More massive discs are expected to be more prone to fragmentation even due to weak perturbations, but are also more likely to reach the limit for self-fragmentation. This study, however, shows that tidal perturbations do not necessarily inhibit fragmentation but are instead capable of inducing fragmentation in disc that otherwise would silently disperse or be accreted by the central star without ever experiencing fragmentation. The perturber star may accrete gas from the target circumstellar disc and may form planets on misaligned or even counter-rotating orbits with respect to the stellar spin. Future work will analyse discs with different masses and more or less massive host stars as well."
    },
    "1005/1005.3151_arXiv.txt": {
        "abstract": "{We present the first far-IR observations of the solar-type stars $\\delta$ Pav, HR 8501, 51 Peg and $\\zeta^2$ Ret, taken within the context of the DUNES {\\it Herschel} Open Time Key  Programme (OTKP). This project uses the PACS and SPIRE instruments with the objective of studying infrared excesses due to exo-Kuiper belts around nearby solar-type stars. The observed 100 $\\mu$m fluxes from $\\delta$ Pav, HR 8501, and 51 Peg agree with the predicted photospheric fluxes, excluding debris disks  brighter than $L_{\\rm dust}/L_\\star \\sim 5 \\times 10^{-7}$ (1 $\\sigma$ level) around those stars. A flattened, disk-like structure with a semi-major axis of $\\sim$ 100 AU in size is detected around $\\zeta^2$ Ret. The resolved structure suggests the presence of an eccentric dust ring, which we interpret as an exo-Kuiper belt with $L_{\\rm dust}/L_\\star$ $\\approx$ 10$^{-5}$.} ",
        "introduction": "The discovery of infrared  excess emission produced by cold, optically thin disks composed of  micron-sized dust grains around main sequence stars is one of the main contributions of IRAS (Aumann et al. \\cite{aumann}). Since the lifetime of  such grains, set by destructive collisions, Poynting-Robertson  drag and radiation pressure, is much  shorter than the ages  of these stars,  one must conclude  that those dust  disks - called debris  disks - are  continuously replenished by  collisions of large rocky bodies (Backman \\& Paresce \\cite{backman}). Observations  of debris  disks provide  powerful  diagnostics from  which  to  learn  about  the dust  content,  its  properties and its spatial  distribution; in addition, since dust sensitively  responds to the gravity of  planets, it can be used as  a tracer of the  presence of planets.   Thus, observations of debris disks around stars of different  masses and ages inform  us about the formation and evolution of planetary  systems, since they are a direct proof  for the  existence of  planetesimals and  an indirect tracer of the presence of planets around stars.  In the Solar  System, the asteroid and Kuiper  belts are examples of debris   systems;  in particular, the Kuiper belt has an estimated   dust luminosity   $L_{\\rm dust}/L_\\star \\sim  10^{-7}$-$10^{-6}$ (Stern \\cite{stern}). IRAS was  only able to detect bright disks, $L_{\\rm dust}/L_\\star >  10^{-4}$, mainly around A and  F stars; ISO extended our  knowledge to a wider  spectral type range  and found a general decline  with the stellar  age (Habing et  al \\cite{habing}, Decin  et al.  \\cite{decin}).  A  remarkable step  forward  has been achieved  by {\\it Spitzer},  studying  debris disks  as  faint as  $L_{\\rm dust}/L_\\star$ several times $10^{-6}$,  their incidence from A to M type  stars, the age  distribution, the presence of  planets, etc. (e.g.  Su  et al.  \\cite{su},  Trilling  et al.   \\cite{trilling08}, Bryden  et  al.  \\cite{bryden09}).   {\\it Spitzer}  has, however,  several limitations.   Its   poor  spatial   resolution   prevent  us from constraining  fundamental disk parameters which require resolved imaging, and the confusion limit inherent in its large beam  limits its detection capability  to cold disks  brighter than the Kuiper  belt by around  two  orders of  magnitude. Also, {\\it Spitzer} is not sensitive longward of 70 $\\mu$m, wavelengths particularly important for the cold disks generally found around Sun-like stars. The far-infrared 3.5 m diameter {\\it Herschel} space telescope (Pilbratt  et  al.    \\cite{pilbratt})  with  its  instruments  PACS (Poglistsch  et al.   \\cite{poglitsch})  and SPIRE  (Griffin et  al. \\cite{griffin})   overcomes    these   limitations,   offering   the possibility  of characterising cold,  $\\sim$ 30  K, debris  disks as faint  as $L_{\\rm  dust}/L_\\star$ few  times $10^{-7}$  with spatial resolution  $\\sim$ 30  AU at  10 pc,  i.e., true  extra-solar Kuiper belts.  DUNES\\footnote{{\\em   DU}st   around   {\\em   NE}arby   {\\em   S}tars, http://www.mpia-hd.mpg.de/DUNES/.}  is a {\\it Herschel} OTKP  designed  to detect  and characterize  cold, faint, debris disks,  i.e., extra-solar analogues to the Kuiper belt, around a statistically  significant sample of  main-sequence FGK  nearby stars, taking advantage of the unique  capabilities of {\\it Herschel} with PACS and SPIRE. The data will  be analysed with radiative, collisional and dynamical dust disk models. A complete description of DUNES goals,  target selection, and  stellar properties will  be presented  elsewhere (Eiroa  et al.   2010,  in preparation). The objectives of the DUNES survey are complementary to those of the  OTKP DEBRIS (Matthews et al.   \\cite{matthews}). Both projects have complementary  star samples, sharing partly some  sources and the corresponding data.  The DUNES  objectives  require  the detection of  very faint excesses at the mJy level, comparable  to  the   photospheric  emission and only a few times the measurement uncertainties. The primary observing  strategy is designed  to integrate  for as long as needed  to detect the 100 $\\mu$m  photospheric flux, subject only to confusion noise limitations.  In this letter  we present  our first results   obtained   during  the   science   demonstration  phase (SDP)  of {\\it Herschel}. Four solar-type G stars were  observed:  $\\zeta^2$  Ret,  $\\delta$  Pav, HR 8501, and 51  Peg.  We also observed q$^1$ Eri as  a test object with a well-known, bright debris  disk;  the  q$^1$   Eri  results  are  presented  in  an accompanying letter (Liseau et al. \\cite{liseau}). Excluding  $\\delta$  Pav, the rest of the stars are shared targets with DEBRIS.  \\section {Observations and data reduction} \\begin{table} \\label{AORs} \\caption{Summary of the SDP DUNES observations.} \\begin{tabular}{lclrcr} \\hline Star          & Obs. ID   &Mode\\tablefootmark{1}& \\multicolumn{1}{c}{Bands} & Scan & \\multicolumn{1}{c}{OT} \\\\ &           &    & \\multicolumn{1}{c}{($\\mu$m)} & Direction    & (sec)     \\\\ \\hline q$^1$ Eri     & 1342187139/40& SM & 100/160 &63$\\degr$/117$\\degr$    &2536       \\\\ q$^1$ Eri     & 1342187141& PS & 100/160 &              &4714       \\\\ q$^1$ Eri     & 1342187142& PS & 70/160  &              &789        \\\\ $\\zeta^2$ Ret & 1342183660& PS & 100/160 &              &1572       \\\\ $\\zeta^2$ Ret & 1342191102/03& SM\\tablefootmark{2} & 100/160 &117$\\degr$/63$\\degr$       &4510       \\\\ $\\zeta^2$ Ret & 1342191104/05& SM\\tablefootmark{2} & 70/160  &63$\\degr$/17$\\degr$ &4510       \\\\ $\\delta$  Pav & 1342187075/06& SM & 100/160 &45$\\degr$/135$\\degr$    &3834       \\\\ HR 8501     & 1342187145/46& SM & 100/160 &63$\\degr$/117$\\degr$    &1890        \\\\ 51 Peg        & 1342187255& PS & 100/160 &              &1731       \\\\ \\hline \\end{tabular} \\tablefoottext{1}{SM = scan map; PS = chop-nod/point-source mode}  \\tablefoottext{2}{Routine phase DUNES observing time (not SDP)} \\end{table}%    The  stars were observed  with PACS  at 70  $\\mu$m (blue),  100 $\\mu$m (green), and 160 $\\mu$m (red). Two observing modes were used - chop-nod/point-source (PS hereafter) and scan map (SM hereafter). Our data set allows us to make  a direct comparison of both modes,  specially in the cases of  q$^1$  Eri  and  $\\zeta^2$  Ret. A  critical  evaluation  of  this comparison will be the subject of a technical note. SM  observations of $\\zeta^2$ Ret  were carried out as DUNES routine  observations. Table  1  gives  some   details  of  the  observations  including  the identification number,  the observing  mode, the wavelength bands, the scan direction angles, and the total duration of the observation (OT).  Data  reduction   was  carried  out   using  the  {\\it Herschel} Interactive Processing Environment (HIPE), version v2.0.0 RC3, and the pipeline script delivered at the December 2009 {\\it Herschel} data reduction workshop,  ESAC, Madrid,  Spain.  The script  provides  all the  tools to  convert pure  raw PACS/{\\it Herschel} data  to flux calibrated and position rectified images.  While the instrument pixel size is 3$\\farcs$2 for the blue and green bands and 6$\\farcs$4 for the red band, the resolution of the final images is set to 1$\\arcsec$/pixel and 2$\\arcsec$/pixel for the blue/green and red bands, respectively. ",
        "conclusions": "Our results show the capability of {\\it Herschel}/PACS to detect and resolve cold dust disks with a luminosity close to the solar Kuiper belt, which will allows us to deepen our understanding of planetary systems, in particular those associated with mature stars. Specifically, our data exclude the existence of cold debris disks with $L_{\\rm dust}/L_\\star \\gtrsim 5\\times10^{-7}$ (1$\\sigma$) around the solar-type stars $\\delta$ Pav, HR 8501 and 51 Peg. On the other hand, the data show that $\\zeta^2$ Ret is a good example where cold disks around nearby stars, very much alike the solar Kuiper belt, can be resolved and studied in great detail with the {\\it Herschel} space observatory."
    },
    "1005/1005.1154_arXiv.txt": {
        "abstract": "{ We use deep observations obtained with the Photodetector Array Camera \\& Spectrometer (PACS) onboard the \\textit{Herschel} space observatory to study the far-infrared (FIR) properties of submillimeter and optically faint radio galaxies (SMGs and OFRGs). From literature we compiled a sample of 35 securely identified SMGs and nine OFRGs located in the GOODS-N and the A2218 fields. This sample is cross-matched with our PACS 100~$\\mu$m and 160~$\\mu$m multi-wavelength catalogs based on sources-extraction using prior detections at 24 $\\mu$m. About half of the galaxies in our sample are detected in at least the PACS 160 $\\mu$m bandpass. The dust temperatures and the infrared luminosities of our galaxies are derived by fitting their PACS and SCUBA 850~$\\mu$m (only the upper limits for the OFRGs) flux densities with a single modified ($\\beta=1.5$) black body function. The median dust temperature of our SMG sample is $T_{dust}=36\\pm8$~K while for our OFRG sample it is $T_{dust}=47\\pm3$~K. For both samples, median dust temperatures derived from \\textit{Herschel} data agree well with previous estimates. In particular, Chapman et al. (2005) found a dust temperature of $T_{dust}=36\\pm7$~K  for a large sample of SMGs assuming the validity of the FIR/radio correlation (i.e., $q=log_{10}(L_{FIR}[{\\rm W}]/L_{1.4\\,{\\rm GHz}}[{\\rm W\\,Hz^{-1}}]/3.75\\times10^{12})$). The agreement between our studies confirms that the local FIR/radio correlation effectively holds at high redshift even though we find $\\langle q\\rangle=2.17\\pm0.19$, a slightly lower value than that observed in local systems. The median infrared luminosities of SMGs and OFRGs are $4.6\\times10^{12}\\,L_{\\odot}$ and $2.6\\times10^{12}\\,L_{\\odot}$, respectively. We note that for both samples the infrared luminosity estimates from the radio part of the spectral energy distribution (SED) are accurate, while estimates from the mid-IR are considerably ($\\thicksim\\times3$) more uncertain. Our observations confirm the remarkably high luminosities of SMGs and thus imply median star-formation rates of 960~M$_{\\odot}\\,$yr$^{-1}$ for SMGs with $S(850\\,\\mu{\\rm m})>5\\,$mJy and 460~M$_{\\odot}\\,$yr$^{-1}$ for SMGs with $S(850\\,\\mu{\\rm m})>2\\,$mJy, assuming a Chabrier IMF and no dominant AGN contribution to the far-infrared luminosity. } ",
        "introduction": "\\indent{ \\textit{Herschel} observations probe the rest-frame far-infrared emission of high-redshift galaxies. Thus, they provide for the first time robust estimates of the infrared luminosities of these high-redshift galaxies and test previous measurements that were based on extrapolation from shorter or longer wavelengths. We here  focus on two populations of high-redshift star-forming galaxies selected at submillimeter (submm) and radio wavelengths. \\\\} \\indent{ Since their discovery in the late 1990s, submillimeter galaxies (SMGs) have become the selection of choice for the most luminous tail of the high-redshift star-forming galaxy population. It has been found that SMGs have a typical redshift of $z\\thicksim2$ \\citep{chapman_2005,pope_2006}, are massive systems \\citep[$M_{\\star}\\thicksim10^{10}-10^{11}\\,M_{\\odot}$,][]{swinbank_2004,tacconi_2006} and are compact \\citep[e.g.,][]{tacconi_2008}. Interferometric observations of their CO molecular gas suggest that the most luminous SMGs ($S_{850\\,\\mu {\\rm m}}> 5$ mJy) are merging sytems \\citep{tacconi_2006,tacconi_2008} with high star-formation efficiencies compared to typical galaxies of a similar mass \\citep{daddi_2008}. Therefore, these SMGs are thought to exhibit very intense (SFR$\\thicksim$1000 M$_{\\odot}$ yr$^{-1}$) short-lived star-formation bursts triggered by mergers and to be the high-redshift progenitors of local massive early-type galaxies \\citep{daddi_2007,daddi_2007a,tacconi_2008,cimatti_2008}. \\\\} \\indent{ Although SMGs provide a powerful tool to constrain the formation and the evolution of high-redshift dusty star-forming galaxies, their selection is subject to strong biases. In particular, because submm observations probe the blackbody emission of dust in the Rayleigh-Jeans regime, they are strongly anti-correlated with the dust temperature ($S_{850}\\propto T^{-3.5}_{dust}$). For a given infrared luminosity, galaxies with hot dust might fall below the detection limit of current submm instruments. The first observational evidence of a missing population of high-redshift dusty star-forming galaxies with hot dust has been  given by \\citet{chapman_2004} using a selection of radio-detected but submm-faint galaxies with UV spectra consistent with high-redshift starbursts. These optically faint-radio galaxies (OFRGs) have a comoving volume density (i.e., $\\thicksim10^{-5}\\,$Mpc$^{-3}$ at $1<z<3$, Chapman et al. \\citeyear{chapman_2004}), stellar masses and sizes comparable to SMGs, and some have a dust temperature of $\\thicksim52~$K \\citep{casey_2009a,casey_2009b}. \\\\} \\indent{ While SMGs and OFRGs are an important component of the high-redshift massive galaxy population, many of their fundamental properties still rely on indirect measurements. In particular, direct determinations of SMG dust temperatures were limited \\citep{kovacs_2006} because they were not done using rest-frame far-infrared observations on both sides of the peak of the SEDs. More importantly, their infrared luminosites are still debated. Indeed, because theoretical simulations of galaxy evolution have had great difficulties to account for the current inferred luminosities/star-formation rates and number counts \\citep{baugh_2005,dave_2010}, they still question whether these luminosities have been overestimated or whether the IMF is significantly more top-heavy than in the local Universe. \\\\} \\indent{ Using deep observations by the Photodetector Array Camera \\& Spectrometer (PACS; Poglitsch et al. \\citeyear{poglitsch_2010}) onboard the \\textit{Herschel} space observatory (Pilbratt et al. \\citeyear{pilbratt_2010}) we will have for the first time robust estimates of the dust temperatures and the infrared luminosities of SMGs and OFRGs. Throughout the paper we use a cosmology with $H_{0}=70 \\rm{km\\,s^{-1}\\,Mpc^{-1}}$, $\\Omega_{\\Lambda}=0.7$ and $\\Omega_{\\rm M}=0.3$. \\\\} \\begin{figure*} \\centering \\includegraphics[width=12.cm]{14616fg1.eps} \\caption{\\label{fig: fit BB}\\small{ Spectral energy distribution (SED) of one SMG (\\textit{left}) and one OFRG (\\textit{right}). Red diamonds present our PACS measurements while green squares present the multi-wavelength ancillary data taken from the literature \\citep{pope_2006,casey_2009a,casey_2009b}. The modified blackbody emission ($\\beta=1.5$) best-fitting the data are shown by solid lines. Dashed lines present the  Dale \\& Helou SED template best-fitting the mid- to far-infrared observations. In the left panel, the blue dotted-dashed line shows the Dale \\& Helou SED template best-fitting the submm and radio photometries. The inset in each panel shows $\\chi^{2}$ vs $T_{dust}$. }} \\end{figure*} ",
        "conclusions": "\\indent{ Figure \\ref{fig: t vs l} depicts the locations of our SMGs and OFRGs on the $T_{dust}$-$L_{IR}$ plane and Fig \\ref{fig: s850 vs lir} shows their locations on the $S_{850\\,\\mu{\\rm m}}$-$L_{IR}$ plane. As already mentioned, OFRGs are biased towards hot dust temperatures; their median $T_{dust}$ is of $47\\pm3$~K and their median $L_{IR}$ is $2.6\\times10^{12}\\,L_{\\odot}$. In contrast, SMGs have lower dust temperatures with median $T_{dust}=36\\pm8$~K and $L_{IR}=4.6\\times10^{12}\\,L_{\\odot}$. We note that lensed-SMGs from A2218 and with $L_{IR}<2\\times10^{12}\\,L_{\\odot}$ exhibit intermediate dust properties and are less biased towards cold dust temperatures than the entire SMG sample. This is because these galaxies would have escaped both the SMG and OFRG selection method without magnification. We also note that \\textit{bright} SMGs (i.e. $S_{850\\,\\mu{\\rm m}}>5\\,{\\rm mJy}$) have higher median infrared luminosities ($L_{IR}=9.6\\times10^{12}\\,L_{\\odot}$) and higher median dust temperatures ($T_{dust}=38$ K) than the entire SMG sample because there is a correlation between $S_{850\\,\\mu{\\rm m}}$ and $L_{IR}$ (Fig \\ref{fig: s850 vs lir}). These estimates are the first direct observational measurements of the dust temperatures and the infrared luminosities of SMGs and OFRGs. \\\\} \\indent{ Our observations reveal that high redshift dusty star-forming galaxies exhibit a wide range of dust temperatures. In particular at low infrared luminosities ($L_{IR}<4\\times10^{12}L_{\\odot}$) the dust temperature dispersion observed in our sample might suggest a higher $T_{dust}$-$L_{IR}$ scatter than that observed by \\citet{chapman_2003} at $z\\thicksim0$. Nevertheless, this conclusion is most likely driven by selection effects because a significant fraction of the galaxies with intermediate dust properties were probably missed by our current sample. Indeed, we note that studying a $L_{IR}$-selected sample of galaxies observed with \\textit{Herschel}, Hwang et al. (in prep) find modest changes in the $T_{dust}$-$L_{IR}$ relation as function of the redshift :  at $z>0.5$, galaxies with $L_{IR}>5\\times10^{10}L_{\\odot}$ are slightly colder ($\\thicksim3\\,$K) than local ones and the scatter of the $T_{dust}$-$L_{IR}$ relation slightly increase at high redshift. \\\\} \\indent{ Though previous estimates of the dust temperatures of SMGs and OFRGs relied on indirect observations, they agree relatively well with our measurements. In particular, \\citet{chapman_2005} found a dust temperature of $T_{dust}=36\\pm7$~K for a large sample of SMGs assuming the validity of the FIR/radio correlation. In order to establish this agreement on a common sample, we applied the same method as \\citet{chapman_2005} to our SMG sample, i.e. we fitted the radio and 850~$\\mu$m photometries with dust SED templates from \\citet{dale_2002} and then translated them into $T_{dust}$ using their R(60,100) to $T_{dust}$ map. With this method we found higher $T_{dust}$ (by $\\thicksim4\\,$K) and $L_{IR}$ ($\\thicksim\\times1.5$ times) than what we obtained using our blackbody analysis. These discrepancies arise because the Dale \\& Helou SED templates assume a FIR/radio correlation with $\\langle q\\rangle=2.34$ \\citep{yun_2001}, while in our samples we find $\\langle q\\rangle=2.17\\pm0.19$ (see also Ivison et al. \\citeyear{ivison_2010b}, this issue). Although this value of $\\langle q\\rangle$ is still in line with results from local systems (which have a dispersion of 0.19~dex), it is also consistent with an evolution of $\\langle q\\rangle$ proportional to $(1+z)^{-0.15\\pm0.03}$ as found by \\citet{ivison_2010}. \\\\} \\indent{ Our results reveal that one can obtain a very reliable estimate of the infrared luminosity of a given galaxy from its radio flux density ($\\sigma[log_{10}(L_{IR}^{Radio}/L_{IR}^{Blackboby})]\\thicksim0.18$ dex) using $\\langle q\\rangle=2.17$. In contrast, the use of the 24 $\\mu$m emission and of the Chary \\& Elbaz SED library yields an inaccurate estimate of the infrared luminosity characterized by a large scatter ($\\sigma[log_{10}(L_{IR}^{24\\,\\mu m}/L_{IR}^{blackbody})]\\thicksim0.48$~dex) and a systematic overestimation ($\\thicksim\\times2$ times) of the most luminous galaxies ($L_{IR}>4\\times10^{12}L_{\\odot}$). We thus find that for the very luminous infrared galaxies studied here, luminosity extrapolations based on the radio emission are considerably more reliable than those based on the mid-infrared emission (see also Elbaz et al. \\citeyear{elbaz_2010} and Nordon et al. \\citeyear{nordon_2010}). \\\\} \\indent{ Using $\\langle q\\rangle=2.17$, we can predict the infrared luminosities of our PACS undetected SMGs. Then, using the Dale \\& Helou SED templates normalized to these infrared luminosities, we can fit their 850~$\\mu$m photometries. We find that for 13 out of 16 undetected SMGs, PACS flux densities inferred using these fits are below the detection threshold of our observations. Of the three sources with PACS fluxes predictions above our detection threshold, one is known to be contaminated by an AGN and the other two are suspected to have wrong redshift estimates \\citep{daddi_2009a}. The PACS nondetections are thus fully consistent with the properties inferred from the detections. This analysis cannot be performed on our PACS undetected OFRGs because they are also undetected at the SCUBA wavelength. \\\\ \\\\} \\indent{ Our observations unambiguously confirm the remarkably large infrared luminosities of \\textit{bright} SMGs (i.e. $S_{850\\,\\mu{\\rm m}}>5\\,{\\rm mJy}$) which correspond to SFRs of $960\\ {\\rm M}_{\\odot}~ {\\rm yr}^{-1}$ (SFR~$[{\\rm M}_{\\odot}~ {\\rm yr}^{-1}] = 1\\times 10^{-10} L_{\\rm IR}~[{\\rm L}_{\\odot}]$, assuming a Chabrier IMF and no dominant AGN contribution to the far-infrared luminosity). Such high SFRs are difficult to reconcile with secular evolution \\citep[e.g.][]{dave_2010} and could correspond to a merger-driven stage in the evolution of these galaxies. This hypothesis is supported by CO observations of \\textit{bright} SMGs which have revealed large CO line-widths and disturbed gas morphologies \\citep{tacconi_2008}. Our observations also confirm that OFRGs exhibit higher dust temperatures than \\textit{faint} SMGs ($S_{850\\,\\mu{\\rm m}}<5\\,{\\rm mJy}$) observed at the same redshift and with equivalent infrared luminosities. While the relatively low median SFRs of OFRGs and \\textit{faint} SMGs ($260\\ {\\rm M}_{\\odot}~ {\\rm yr}^{-1}$ and $106\\ {\\rm M}_{\\odot}~ {\\rm yr}^{-1}$ respectively) could be explained by secular evolution, we need to understand why they exhibit different dust temperatures and to study a possible link with \t\\textit{bright} SMGs. A clear evolutionary picture will require detailed studies of the dust and molecular gas distribution in a sample of high-redshift star-forming galaxies unbiased towards any particular dust temperature. This SFR-selected sample can now be built with our ongoing deep \\textit{Herschel} observations. \\\\} \\acknowledgement{ PACS has been developed by a consortium of institutes led by MPE (Germany) and including UVIE (Austria); KU Leuven, CSL, IMEC (Belgium); CEA, LAM (France); MPIA (Germany); INAF-IFSI/OAA/OAP/OAT, LENS, SISSA (Italy); IAC (Spain). This development has been supported by the funding agencies BMVIT (Austria), ESA-PRODEX (Belgium), CEA/CNES (France), DLR (Germany), ASI/INAF (Italy), and CICYT/MCYT (Spain). }"
    },
    "1005/1005.3684_arXiv.txt": {
        "abstract": "{% Three-dimensional numerical simulations of magnetoacoustic wave propagation are performed in a sunspot atmosphere with a computational domain covering from the photosphere to the chromosphere. The wave source, with properties resembling the solar spectrum, is located at different distances from the axis of the sunspot for each simulation. These results are compared with the theory of mode transformation and also with observational features. Simulations show that the dominant oscillation frequency in the chromosphere decreases with the radial distance from the sunspot axis. The energy flux of the different wave modes involved, including de Alfv\\'en mode, is evaluated and discussed.} ",
        "introduction": "Observations of waves in sunspots show a variety of behaviours depending on the height and the region observed. In the umbral photosphere their power spectrum has a peak at 3.3 mHz (\\hbox{period} of 5 min) and their \\hbox{amplitudes} are a\\-round a hundred m s$^{-1}$. The frequency of the power peak varies with the height, and at the chromosphere the oscillatory pattern is dominated by waves in the 5-6 mHz band (period of 3 min) with amplitudes of several kilometers per second which \\hbox{present} a saw-tooth profile. Several \\hbox{mechanisms} have been proposed to explain the chromospheric three min\\-ute oscillations: a resonant chromospheric cavity \\citep{Zhugzhda+etal1985}; non-linear interaction of pho\\-to\\-spheric modes \\citep{Gurman+Leibacher1984}; and slow magneto-acoustic \\hbox{mode} \\hbox{field} \\hbox{aligned} propagation from the photosphere to the chromosphere in the 5-6 mHz band \\citep{Centeno+etal2006}.  The interpretation of the observed oscillations in terms of MHD waves possess several difficulties. The velocity of propagation of the magneto-acoustic modes depends on the acoustic velocity ($c_S$) and Alfv\\'en velocity ($v_A$), both stratified with height. The magnetized atmosphere of sunspots changes from being gas dominated in the photosphere and below, ($cs>v_A$) to field dominated at the high chromosphere ($c_S<v_A$). Around the layer where $c_S=v_A$ the phase speeds of all modes are similar and different waves can interact, which produces the mode trans\\-for\\-ma\\-tion \\citep{Bogdan+etal2003, Cally2006, Khomenko+Collados2006}. The direction and efficiency of the mode transformation depends on the frequency of the wave and the angle between the wave vector and the magnetic field \\citep{Cally2006}.  Most of the simulations of mode refraction,  transformation, etc. have been developed for waves with high frequencies, $i.e.$ above the cut-off frequency of the solar atmosphere, and using 2D atmospheres. We aim to study these phenomena through numerical simulations of waves with the periods observed and using realistic atmospheres of sun\\-spots in a 3D case. One of the most important benefits of 3D simulations of wave propagation in magnetized atmospheres is the appearance of the Alfv\\'en mode. The work by \\citet{Felipe+etal2010} points out that the conversion from the slow to the Alfv\\'en mode is effective when the wave vector forms a certain angle different from zero with the magnetic field. Conversion to Alfv\\'en waves only was produced when the driver was located out of the axis of the sunspot, where the magnetic field has some inclination. In this contribution, we explore further the conclusions of that work to study waves with 3-5 min periodicity found in sunspots, and analyze several simulations with the driver located at different distances from the center of the sunspot.  \\begin{figure*} \\includegraphics[width=165mm]{vs_x5_t150_expandz.eps} \\caption{Variations of longitudinal velocity scaled by a factor $\\rho^{1/2}$ at an elapsed time t=820 s after the beginning of the simulations with the driver force at 5 Mm from the axis of the sunspot. Grey inclined lines are magnetic field lines.} \\label{fig:3D_long} \\end{figure*}   \\begin{figure*} \\includegraphics[width=165mm]{vf_x5_t150_expandz.eps} \\caption{Variations of transversal velocity scaled by a factor $\\rho^{1/2}$ at an elapsed time t=820 s after the beginning of the simulations with the driver force at 5 Mm from the axis of the sunspot. Grey inclined lines are magnetic field lines.} \\label{fig:3D_trans} \\end{figure*} ",
        "conclusions": "\\begin{itemize}  \\item The mechanism that produces the change in frequency of oscillations in the umbra from the photosphere to the chromosphere is the linear propagation of waves with 3 minute power which come directly from the photosphere.  \\item The dominant wave frequency in the chromosphere decreases with the radial distance, due to the reduction of the cut-off frequency far from the axis.  \\item Only waves in the 5-6 mHz frequency band can provide energy to the chromosphere through slow acoustic mode, since waves with lower frequencies form evanescent waves and do not carry energy, and the fast magnetic mode is reflected back towards the photosphere.  \\item The conversion to the Alfv\\'en mode is effective where the magnetic field has some inclination, but its energy flux is small compared to the acoustic energy flux.  \\item Most of the energy is confined to the region below the height where the sound and Alfv\\'en velocity are similar.  \\end{itemize}"
    },
    "1005/1005.1648_arXiv.txt": {
        "abstract": "We present the first non-LTE, co-moving frame molecular line calculations of a star-forming cluster simulated using smoothed particle hydrodynamics (SPH), from which we derive high-resolution synthetic observations. We have resampled a particle representation onto an adaptive mesh and self-consistently solved the equations of statistical equilibrium in the co-moving frame, using \\textsc{torus}, a three-dimensional adaptive mesh refined (AMR) radiative transfer (RT) code. We verified the applicability of the code to the conditions of the SPH simulation by testing its output against other codes. We find that the level populations obtained for optically thick and thin scenarios closely match the ensemble average of the other codes. We have used the code to obtain non-LTE level populations of multiple molecular species throughout the cluster and have created three-dimensional velocity-resolved spatial maps of the emergent intensity. Line profiles of cores traced by \\nnh~(1-0) are compared to probes of low density gas, \\low~(1-0) and \\high~(1-0), surrounding the cores along the line of sight. The relative differences of the line-centre velocities are shown to be small compared to the velocity dispersion, matching recent observations. We conclude that one cannot reject competitive accretion as a viable theory of star formation based on observed velocity profiles. ",
        "introduction": "\\label{Intro} Large-scale hydrodynamic collapse calculations of turbulent clouds (e.g. \\citealt*{BaBoBromm03,KruKleMcK07a}) demonstrate the fragmentation of molecular clouds into clusters which in turn typically fragment into high-density cores and subsequently fragment into protostars. These simulations predict the statistical properties of the distribution of young stellar objects such as the initial mass function, the binary fraction and the star formation rate (e.g. \\citealt{Zin84,BoBaVi03,BaBoBromm03,Bate09a}). It is expected that the fate of these objects is strongly affected by their evolution history, specifically the competitive accretion of gas as they move through their environs. Star formation in isolation \\citep*{AdamsLadaShu87} necessarily neglects the dynamic interactions between these objects, which can cause them to be ejected at great velocities from their nascent gas-rich core, truncating their accretion \\citep*{Klessen98,BaBoBromm02b} and it will be possible, using optically thin molecular tracers, to observe any motion of the object relative to the enveloping gas \\citep*{Walshetal04}.  Comparisons of the collapse calculations with observations have generally relied on statistical tests of protostellar properties such as the IMF and binary fractions (\\citealt{Bate09b}, but see also \\citealt{Kurosawaetal04}). Although undoubtedly useful, such comparisons neglect the information encoded in the geometrical and dynamical properties of the star-forming gas and considering the wealth of observational data available (e.g. \\citealt{NarHeyBru08,BucCurRob10}) these tests should provide a further, robust appraisal of the models. In order to make comparisons between theory and observation, given the intractability of the inverse problem, it is necessary to transform the temperature and density characteristics of a simulated cloud into a synthetic intensity map comparable with those obtained by observation. That being so, this paper aims to improve upon the analysis of the relative motions of cores to their nascent envelopes in the cluster models of \\citet{BaBoBromm03} by \\citet{Ayliffeetal07}. Both \\cite{Ayliffeetal07} and \\cite{Walshetal04} look for evidence of supersonic ballistic motions of cores relative to their inchoate envelopes. Assuming that the velocities of the core and the envelope are in the line of sight it is possible to determine the difference between them.  By performing a non-LTE  radiative transfer calculation of the molecular level populations within the cluster using \\textsc{torus} (see \\citealt{Harries00, Harries04, Symington05a} for details), we seek to improve upon the comparison between simulation and observation carried out by \\cite{Ayliffeetal07} and reconcile the differences between observation and theory.  In order to understand the chemical and physical conditions in these structures it is necessary to understand the role that molecules play in radiation transfer in these cool, dense regions. For over 40 years (e.g. \\citealt{Dieter64, WilsonPenzias70}) molecular line radiation has regularly been used to trace key physical parameters such as density, temperature and velocity. Today's (sub-)millimetre observatories enable the observation of some of the coolest, most chemically rich regions of the galaxy. Furthermore, the Atacama Large Millimeter Array (ALMA) promises unprecedented resolution of these objects therefore it is vital that the interpretation of the data is sound, necessitating the use of non-symmetric models incorporating complex physics. This was first recognized by \\citet{Bernes79} and today many line radiative transfer codes exist (e.g. \\citealt{Juvela97, Wiesemeyer97, Rawlings&Yates01}). Until relatively recently, only spherically or axially symmetric core collapse simulations of line transfer had been performed (e.g. \\citealt{Tsamis08,Pavetal08}) however, due to recent advances in computational power, it is now possible to compute three-dimensional line transfer for these complex density structures in star-forming regions. Here we use the \\torus code, since the use of adaptive mesh refinement (AMR) enables us to treat the huge dynamic range in density and linear scale associated with the SPH cluster calculations.  In section \\ref{Implementation} we present our implementation of a non-LTE molecular line radiative transfer module for \\textsc{torus}. In section \\ref{benchmarking} we demonstrate agreement with other codes in some of the benchmark problems set out in \\citet{vZ02} and verify the accuracy of our raytracing routine for generating datacubes. In section \\ref{sphtogrid}, we present an efficient method for mapping irregular SPH data on to an AMR grid and in section \\ref{sph} we present the results of its application to the hydrodynamic clustered star-formation simulation of \\cite{BaBoBromm03}. Finally, we discuss the results of our analysis of the cluster simulation and conclude on our findings and on the importance of radiative transfer (RT) analysis of star formation as a whole. ",
        "conclusions": "\\label{conclusion} This paper demonstrates the molecular line radiative transfer capabilities of \\textsc{torus}, a radiative transfer code that uses AMR to resolve the fine details of complex three-dimensional environments. The molecular line transfer module was compared against other codes in a scenario where LTE conditions did not apply and the results showed excellent agreement.  In environments such as molecular clouds, gas that is dense enough to emit with sufficient intensity in a given molecular line may only exist in small regions when compared with the size of the whole cloud. Consequently, fine discretization is only required over a small volume. Thus, we devised an efficient method of mapping irregular SPH data onto an AMR grid with minimal discretization error. This technique permitted the study of the expected molecular line emission from a hydrodynamical simulation of molecular cloud collapse using SPH with our grid-based radiative transfer code. Having obtained non-LTE level populations throughout the cluster for tracers of high-- and low--density gas over 5 timesteps covering the creation of 50 protostellar objects we were able to produce and analyse synthetic observations. We have shown in this paper that clusters exhibiting competitive accretion are able to reproduce the properties of relative core motions found in observation.  As well as enabling us to accurately model line profiles, the mapping of density distributions from the particle-based representation onto an AMR grid has important implications for coupling radiation transfer and hydrodynamics. For example, using the technique outlined in section \\ref{sphtogrid}, it is now possible to conduct a full, polychromatic multiple scattering treatment of the transport of radiation in an SPH simulated circumstellar disc using high-accuracy grid-based techniques \\citep{AcrHarRun10} without having to resort to the flux-limited diffusion approximation employed in earlier radiative transfer simulations (e.g. \\citealt{WhiteBate04,KruKleMcK07a,Bate09b}). However, extending such calculations to a cluster collapse model is currently not tractable given the very large dynamic range of linear scales necessary to accurately treat the radiation transport of optically thin-to-thick boundaries, but, with further advances in computational power this kind of calculation will become a reality, bringing with it enhanced certainty in the conclusions that can be reached using these simulations."
    },
    "1005/1005.1681_arXiv.txt": {
        "abstract": "We examine the effect of mass and force resolution on a specific star formation (SF) recipe using a set of \\emph{N}-body/Smooth Particle Hydrodynamic simulations of isolated galaxies. Our simulations span halo masses from $10^9$ to $10^{13}$ $M_{\\odot}$, more than four orders of magnitude in mass resolution, and two orders of magnitude in the gravitational softening length, $\\epsilon$, representing the force resolution. We examine the total global star formation rate, the star formation history, and the quantity of stellar feedback and compare the disk structure of the galaxies. Based on our analysis, we recommend using at least $10^4$ particles each for the dark matter and gas component and a force resolution of $\\epsilon \\approx 10^{-3} R_{vir}$ when studying global SF and feedback. When the spatial distribution of stars is important,  the number of gas and dark matter particles must be increased to at least $10^5$ of each. Low mass resolution simulations with fixed softening lengths show particularly weak stellar disks due to two-body heating. While decreasing spatial resolution in low mass resolution simulations limits two-body effects, density and potential gradients cannot be sustained. Regardless of the softening, low-mass resolution simulations contain fewer high density regions where SF may occur. Galaxies of approximately  $10^{10}$ $M_{\\odot}$ display unique sensitivity to both mass and force resolution. This mass of galaxy has a shallow potential and is on the verge of forming a disk. The combination of these factors give this galaxy the potential for strong gas outflows driven by supernova feedback and make it particularly sensitive to any changes to the simulation parameters. ",
        "introduction": "\\label{introsec}  The formation of stars is a highly non-linear process that is dependent on a number of different physical mechanisms, possibly including but not limited to: gravitational \\citep{Zinnecker84, Larson85} or turbulent \\citep{PadoanANDNordlund02, BateANDBonnell04} fragmentation, competitive accretion \\citep{BonnellANDBate06}, magnetic fields \\citep{Shu04}, stellar feedback \\citep{Silk95, AdamsANDFatuzzo96}, and the coalescence of protostars \\citep{MurrayANDLin96}. On larger scales, SF is driven by gravitational interactions between galaxies \\citep{Barton00, Li08}, external gas accretion \\citep{Kauffmann06}, and disk instabilities such as bars \\citep{Ho97, Sheth05} and spiral arms \\citep{Roberts69, Tamburro08}. Furthermore, energy from active galactic nuclei and supernova feedback \\citep{Martin99, Kaviraj07, Bundy08} regulate SF. Hence, SF is determined by the interaction between many large and small scale physical processes.  SF is commonly studied on a galactic scale using \\emph{N}-body/Smoothed Particle Hydrodynamic (SPH) simulations because of their ability to model processes over orders of magnitude in space and mass while automatically resolving denser areas to a higher degree. In fact, it is only very recently that an Eulerian hydrodynamic simulation has followed the formation of a galaxy in a cosmological context to a redshift of zero \\citep{Gibson09}. Cosmological SPH simulations \\citep{Keres05, Governato07, Brooks09} contain one or more galaxies forming within a larger, cosmologically determined dark matter (DM) structure and are crucial for replicating the tidal interactions with other galaxies that may trigger SF \\citep{Okamoto05, Kaufmann06, Governato07}. These cosmological simulations are used for studying the connection between SF and UV background radiation \\citep{Bullock00, Kravtsov04, Moore06}, mergers \\citep{Abadi03a}, tidal and ram pressure stripping \\citep{Mayer06}, and gas accretion by galaxies \\citep{Keres05, Brooks09}. On a smaller scale, simulations of isolated galaxies give a resolved model of SF within a galaxy and its effect on the galactic structure \\citep{Valenzuela03, Debattista04, Tasker06, Stinson07, Roskar08}.  Even in high resolution SPH simulations of the smallest isolated galaxies, the formation of individual stars is unresolved. In order to describe the non-linear, sub-resolution process in SPH simulations, SF is implemented using sub-grid models wherein the complex physical processes are averaged over the scale of the smallest resolution element. Comparisons with Eulerian grid models of SF in the interstellar medium demonstrate the viability of sub-grid SF methods in SPH simulations \\citep{Tasker06}. Given basic information about the surrounding gas, such as density and temperature, statistical models of SF can reproduce most observed global SF properties of galaxies in well-resolved simulations. The question remains, however, what is necessary for a simulation to be ``well-resolved\"? Although the qualitative behavior of SF in SPH simulations has been discussed in previous papers, here we will establish the minimum spatial and mass resolution necessary to model different aspects of SF in galaxies simulated with the widely used code GASOLINE \\citep{Wadsley04} over a wide range of masses. As well as delineating the behavior of SF with GASOLINE as a function of resolution, these results  may be applicable to other SPH simulations in which individual gas particles represent single-phases of the ISM, SF occurs in suitably cold and dense gas particles, and energy from supernova is transferred to nearby gas particles, e.g. \\citet{Katz96, ThackerCouchman01, Keres05, Saitoh08}.   \\citet{Stinson06}, hereafter S06, described a specific sub-grid model for SF and a blastwave model for stellar feedback in GASOLINE, laid out a physical justification for the models, and presented several different tests of the SF and feedback models. To judge the effects of different SF criteria and model parameters, S06 used an isolated Milky Way-mass disk galaxy. When using the best-fit parameters, this model galaxy was able to reproduce the Schmidt law, $\\Sigma_{SFR}\\propto \\Sigma_{gas}^{1.4}$ \\citep{Kennicutt98}, with a near constant star formation rate (SFR). S06 further simulated this isolated galaxy at three different mass resolutions ($2000$, $10 000$ and $50 000$ initial gas particles and a gas disk mass of $10^9  M_{\\odot}$). While fewer stars formed in simulations run at lower resolution, the global SFR started to converge for $10 000$ initial gas particles. \\citet{Brooks07} found that the total amount of SF and the metallicity converged for over a few thousand DM particles per galaxy in cosmological simulations. \\citet{Governato07} further tested this SF recipe when calibrating the SF parameters through comparisons of the SFR and velocity dispersion of an isolated dwarf and a MW-mass galaxy to observations. This SF and feedback recipe, albeit using a different set of parameters, has been successful in reproducing the Tully-Fisher Relation \\citep{Governato07},  the internal structure of disk galaxies \\citep{Governato04, Governato07, Roskar08, Stinson09}, the column density distributions of damped Layman $\\alpha$ systems \\citep{Pontzen08,Pontzen10}, and the observed stellar mass-metallicity relationship \\citep{Brooks07}.  Despite this success, it is possible that SF in galaxy simulations outside of the previously tested range of resolution could behave inconsistently. \\citet{Saitoh08} found that when using a similar SPH/N-body code, in very high resolution simulations ($3.5 \\times 10^2 - 10^3 M_\\odot$ per particle), increasing the threshold density for SF and decreasing the efficiency resulted in more realistic disks which were both thinner and clumpier. If simulations with different resolutions show different properties, such as a range of gas densities, the SF recipe may need to be adjusted at these extremes. As well as affecting the robustness of the particular sub-grid SF model used, resolution also affects the dynamics of galaxy simulations and impacts the SF through the alteration of the density structure of the galaxy \\citep{SpringelANDHernquist02, Governato07, Kaufmann07, Naab07, Foyle08}. Unlike the recipe put forth in \\citet{SpringelANDHernquist03a}, which includes a sub-resolution multiphase model of the ISM, the S06 recipe assumes a single pressure, density and temperature for the entire gas particle. \\citet{SpringelANDHernquist03a} developed their multiphase model in part to minimize resolution effects but at the cost of small scale spatial structure of the multi-phase ISM. Because GASOLINE uses a single phase ISM for each particle it is especially important that the density structure of the ISM be well resolved to correctly model the gas associated with SF. Understanding the impact of resolution becomes particularly important when analyzing large-volume, cosmological simulations that contain a number of galaxies with different masses, each resolved to different degrees. The wide range of galactic masses in these cosmological simulations mean SF must remain accurate in a variety of different galactic environments, from large galaxies with high densities of gas to small ones from which gas is easily expelled. It is, therefore, necessary to establish how SF behaves at different mass and spatial resolutions over a range of galactic masses.  There is a large body of work on the effect of both mass resolution (the mass of the particles and inversely related to the number of particles) and force resolution (the gravitational softening length, $\\epsilon$) on the dynamics of SPH simulated galaxies. The effects of insufficient resolution can be largely grouped as those relating to over-cooling \\citep{HutchingsANDThomas00,SpringelANDHernquist02}, two-body heating \\citep{Steinmetz97, Governato07,Mayer08}, artificial viscosity \\citep{Kaufmann07, Wadsley08}, and small-scale gravitational support \\citep{BateANDBurkert97}. The first two effects primarily relate to the mass resolution whereas the third is connected to the force resolution. These effects can make the galaxy either more or less conducive to SF and they can even impact each other. Understanding how each alone impacts the dynamics of a simulated galaxy, however, is the first step to understanding how together they may influence SF.  The first potential problem for accurately modeling SF in low resolution SPH simulations is that over-cooling of gas in galactic halos simulated with small numbers of particles \\citep{SpringelANDHernquist02} may allow the gas to reach the density threshold for SF too quickly. This cooling, also known as in-shock cooling \\citep{HutchingsANDThomas00} happens when low resolution causes the broadening of shock fronts and contact discontinuities and, as a result, the gas moving through them is heated relatively slowly. This longer heating time-scale gives the gas the opportunity to cool radiatively, which reduces the post-shock cooling time by as much as a factor of two and lowers the maximum shock temperature by up to $50\\%$ in shock-tube simulations \\citep{HutchingsANDThomas00}.  In low mass resolution simulations, two-body heating and the processes related to it may affect SF by changing both the dynamics and the gas physics. Two-body heating results from approximating a smooth potential with discrete masses and causes an artificial increase of kinetic energy. In particular, the presence of  large collisionless bodies, such as dark matter (DM) particles, can heat the disk and destroy galactic disk structure \\citep{LaceyANDOstriker85}. Even when the rotational disk is not fully degraded through two-body heating, numerical relaxation can transfer angular momentum from the disk to the halo \\citep{Kaufmann07}. This results in a weaker disk component \\citep{Governato04} and excessively concentrated disks \\citep{Governato07}. In addition to changing the dynamics, two-body heating increases the gas temperature through artificial viscous dissipation \\citep{Steinmetz97}, making it increasingly difficult for gas to reach temperatures below the temperature threshold for SF.  One final way in which the small number of particles in a simulation may affect SF is through the reduced spatial resolution of the gas. Simulations with small numbers of particles have their gas properties calculated over larger smoothing lengths, $h$, which in turn adds artificial viscosity and dissipation \\citep{Kaufmann07} and results in an increased loss of angular momentum and smaller, more concentrated disks. The smoothing of gas properties over larger volumes also results in fewer regions with sufficiently cool temperatures and high densities for SF.  The ability of regions to collapse to form stars in \\emph{N}-body simulations is likewise sensitive to the force resolution, i.e. the gravitational softening length. Large values of $\\epsilon$ spread the potential over a larger volume, resulting in decreased density fluctuations and producing puffier disks \\citep{Kaufmann07}. Increasing $\\epsilon$, therefore, decreases the amount of gas able to reach sufficient densities for SF. Larger values of $\\epsilon$ also hinder the formation of bars as bars are only formed in the non-axisymmetric modes of the density field that are larger than $\\epsilon$. Without bars, spiral structure is damped and less angular momentum is transferred outward \\citep{Kaufmann07} causing a lower concentration. Without the high densities of gas produced in bars, spiral arms, and in the center of the galaxy, SF may be reduced.  When the Jeans length is comparable to either $\\epsilon$ or $h$ the relationship between them becomes particularly important. When $\\epsilon \\approx h$, gravitational forces on scales smaller than the softening length are reduced, causing artificial support. Conversely, when $\\epsilon < h$ gravitational forces dominate over pressure forces \\citep{BateANDBurkert97}. Reducing $\\epsilon$ relative to $h$ increases the relative gravitational importance, while increasing it results in greater pressure support. Once the Jeans length begins to be unresolved, gas dynamics, and through that SF, become dependent on the relative values of $\\epsilon$ and $h$  \\citep{BateANDBurkert97}. The extent to which relative values of $\\epsilon$ and $h$ affect SF is itself dependent on the implementation of the SF recipe. % For this reason, it is important that the Jeans mass and length be greater than the respective mass and softening lengths of the star forming particles.  As well as the effects of resolution on SF, the effects on the supernova feedback that regulates it must also be understood. All SPH feedback recipes must contend with the feedback energy transferred to the nearby dense gas being radiated away too quickly, resulting in little hydrodynamical effect \\citep{Katz96}. One commonly used solution is to inject some of the SN energy in the form of kinetic rather than thermal energy  to ``wind particles'' \\citep{SpringelANDHernquist03a}. Energy from feedback not used in the wind is necessary to maintain the equation of state in the presence of radiative losses. The effect of mass resolution on the wind produced in this recipe have been examined by \\citet{DallaVecchia08}, among others. A second solution to simulating the sub-grid adiabatic phase of supernova explosions is to temporarily disabling gas cooling. This is the approach taken by the blastwave model, first presented in \\citet{Gerritsen97} and \\citet{ThackerCouchman00} and later developed for Gasoline in S06. As no kinetic energy is explicitly transferred in this recipe, resolution will have a different effect. For example, in low mass resolution simulations using the S06 recipe, supernova events occur less frequently but are more powerful. In such simulations, the energy released per event is greater, that energy is distributed over fewer particles, and the cooling for each particle is disabled for a greater period of time (S06). While this recipe produces the same galaxy-and-time-averaged quantities irrespective of resolution, (for instance, cooling is on average disabled for the same total mass of particles in a high resolution simulation as in a lower resolution one at any given time), it is an open question as to whether the different spatial and temporal distributions of supernova events in low and high mass resolution simulations will produce structurally different galaxies. Furthermore, the high spatial detail produced by using large numbers of particles may allow for the formation of non-spherically symmetric structures, such as galactic winds. Such structures would likewise be affected by resolution induced changes to the artificial viscosity of gas particles \\citep{Kaufmann07} and changes to the disk scale height and the geometry of the disk \\citep{MacLow88, MacLow89}.  In this paper, we explore the effects of force and mass resolution on S06 SF and stellar feedback in different mass galaxies. When analyzing galaxies with this SF recipe, it is necessary to know the lower limit of mass and force resolution above which SF and the stellar disk properties converge. Correctly determining the amount and location of SF is vital in simulations in galaxies because SF 1) produces and distributes metals throughout the galaxy, 2) affects the distribution of matter in the galaxy through feedback, and 3) enables us to relate simulations to observations. Therefore, we ask: what resolutions are sufficient to establish the galactic SF rates, the history of SF, the amount of stellar feedback, and the distribution of stellar and gaseous matter? To this end, we simulated a series of isolated galaxies of different masses at different mass and force resolutions and analyzed properties relating to SF. Isolated galaxies have the advantage of requiring less computational expense while allowing us to separate out environmental effects from those caused by different resolutions. We then analyze the SF and stellar feedback in the models and relate them to the history and structure of the galaxies to determine the resolution necessary for convergence and to describe the effects of resolution.  The outline of our paper is as follows. In \\S 2, we introduce the models and describe the computational methods used. The results of these simulations are described in terms of global SF (\\S 3.1), stellar feedback (\\S 3.2) and stellar distribution (\\S 3.3). We address the physical connection between the SF, feedback, and structure and discuss implications of these results for large volume simulations in \\S 4.  We conclude \\S 5 with a list of recommendations for future cosmological simulations. ",
        "conclusions": "\\label{conclusionsec} With limited computational resources, a compromise is always necessary between sampling a large number of galaxies and resolving those galaxies well. The practical question is not at what level of resolution do all measurements of the galaxy converge, but what minimum resolution is needed to answer a specific question. Studying the bright end of the galactic luminosity function, for example, requires very different resolution than studying the SFHs of dwarf galaxies.  Based on our isolated galaxies, we recommend resolutions for various modeling goals in cosmological SPH simulations for the S06 SF and feedback recipes. These recommendations may also be applicable to any SPH simulation with single-phase gas particles in which supernova feedback energy is transferred to the surrounding gas particles. As these are based on isolated galaxies, they come with a number of caveats when extrapolating to cosmological simulations. Namely, our galaxies have purely spherical accretion, have no mergers, and have no outside tidal forces. These limitations affect the properties of accreted gas as well as the amount of angular momentum transferred to and from the galaxy. Our initial conditions also only sample one standard angular momentum ($\\lambda$ = 0.039) and one mass profile (NFW) and assume the \\citet{Kroupa93} initial mass function. Despite these caveats, our set of simulations covers a wide range of masses, mass resolutions, and force resolution and produce both ``bursting\" and sustained SF at different stages. Together they provide a point of reference for determining the necessary resolution of galaxies. With these limitations in mind, we make the following recommendations.  \\begin{itemize} \\item \\emph{For simulations similar to ours with gas particle masses less than $10^7 M_\\odot$, the Jeans mass of star forming gas is resolved} This implies that our values of $n_{min}$ and $c^*$ in  the SF recipe are suitable for our range of resolutions using our cooling recipe.  \\item \\emph{Fewer than $10^4$ DM particles or initial gas particles or softening and smoothing lengths greater than $10^{-3} R_{vir}$ do not produce accurate global SFRs and supernova feedback.} Over a range of masses, we find that the average global SFR converges for isolated galaxy simulations of $10^4$ particles and for softening lengths smaller than $10^{-3} R_{vir}.$ Gas loss and accretion also converge for similar values (except in the $10^6$ DM particle, $10^{10} M_\\odot$ simulation). While the connection between mass loss and SF is clear, it is remarkable that the mass loss convergence happens at the same values over almost the entire range of galactic masses studied, despite the differences in the galactic potentials.  \\item \\emph{SFHs for galaxies do not converge for simulation using less than $10^4$ DM particles and less than $10^4$ initial gas particles or for simulations with softening lengths larger than  $10^{-3} R_{vir}$. In galaxies near $10^{10} M_\\odot$, the interaction between feedback and the shallow potential wells create spurious effects and prevents convergence at all mass resolutions.} In general, increasing mass resolution results in more peaked bursts of SF followed by a smooth decrease in the SFR. The $10^{10} M_{\\odot}$ galaxy, however, shows a wide range of behavior over the mass resolutions tested with no clear convergence. For this mass of galaxy, raising the number of particles causes the creation of a disk and raising it further results in a sudden increase in the amount of mass lost through supernova feedback. The SFHs of this particularly sensitive galaxy reflect this: at low mass resolution the lack of a coherent disk results in episodic SF; slightly better mass resolution simulations show higher peaked SFRs; the highest resolution galaxy shows less SF as stellar feedback becomes more efficient.  \\item  \\emph{The spatial distribution of stars is not robust for softening lengths of $ \\epsilon \\ge 10^{-2} R_{vir}$ or in simulations with fewer than $10^5$ particles.} Low mass and force resolution affect the spatial distribution of the stars by making the galaxy less disky with a less steep central density profile. Improving the mass resolution reduces the angular momentum loss and shrinking the softening length similarly reduces artificial pressure and allows disks to collapse more fully. Disk properties in galaxies with $10^4$ particles are only starting to converge and $10^{5}$ particles are generally necessary for the results to be robust. \\end{itemize}  For galaxies much smaller than $10^{10} M_{\\odot}$, low temperature cooling from metals \\citep{Mashchenko2008} and $\\mathrm{H_2}$ is necessary for gas to reach temperatures below the virial temperature. This is particularly important when predicting the effect of stellar feedback in dwarf galaxies or at high redshift. High resolution simulations in combination with low temperature cooling and $\\mathrm{H_2}$ formation would better allow us to probe the star formation within dwarf galaxies, including the effects of feedback on rotation curves and episodic versus sustained star formation. In future work, we will examine the effect of $\\mathrm{H_2}$ on cooling in galaxies of different masses and present an $\\mathrm{H_2}$-based SF model"
    },
    "1005/1005.5294_arXiv.txt": {
        "abstract": "Malin\\,1 is a unique, extraordinarily large low surface brightness galaxy. The structure and the origins of the galaxy are poorly understood. The reason for such a situation is an absence of detailed observational data, especially, of high-resolution kinematics. In this Letter we study the stellar kinematics of the inner part ($r \\leq 15$\\,kpc) of Malin\\,1.  We present spectroscopic arguments in favour of a small galaxy -- Malin\\,1B -- being a companion probably interacting with the main galaxy -- Malin\\,1. This object is clearly seen in many published images of Malin\\,1 but is not mentioned in any astronomical databases. Malin\\,1B is located at the projected distance of 14 kpc from the Malin\\,1's nucleus and has small -- 65$\\pm$16 km/s -- relative velocity, which we determined for the first time. We suggest that ongoing interaction with Malin\\,1B can explain main morphological features of the Malin\\,1's central region -- two-armed spiral structure, a bar, and an external one-armed spiral pattern.  We also investigated the large scale environment of Malin\\,1 and postulate that the galaxy SDSS\\,J123708.91+142253.2 might be responsible for the formation of extended low-surface brightness envelope by means of head-on collision with Malin\\,1 (in the framework of collision scenario proposed by Mapelli et al. 2008). To test the collisional origins of Malin\\,1 global structure, more observational data and new numerical models are needed. ",
        "introduction": "Malin\\,1 is one of the most unusual galaxies known to date. The galaxy was accidentally discovered in the course of a systematic survey of the Virgo cluster region designed to detect extremely low-surface brightness (LSB) galaxies (Bothun et al. 1987). According to observational characteristics, Malin\\,1 remains unique among other known galaxies.  Radial extent of its stellar disc in the $V$ passband reaches out $\\approx$60$\\arcsec$ or $\\approx$90 kpc (adopting distance $D_L$=366 Mpc and scale 1.51 kpc/1$\\arcsec$ from the NASA/IPAC Extragalactic Database (NED)) and the disc scale length is about 45$\\arcsec$ or 68 kpc (Bothun et al. 1987). Deep $R$-band data show even larger radial size of the disc -- about $80\\arcsec$ or 120 kpc -- and somewhat smaller scale length value --  33$\\arcsec$ or 50 kpc (Moore \\& Parker 2006). Therefore, Malin\\,1 possesses the largest stellar disc of any known spiral galaxy. The disc is of extremly low surface brightness -- extrapolated central surface brightness is $\\mu_0(V) \\approx 25\\fm5$/$\\square\\arcsec$ (Bothun et al. 1987), $\\mu_0(R) \\approx 24\\fm7$/$\\square\\arcsec$ (Moore \\& Parker 2006). Malin\\,1 has a very low brightness, but due to its enormous size the galaxy's total optical luminosity is high: $M_V \\approx -22\\fm9$ (Pickering et al. 1997). Based on these characteristics, Malin\\,1 is often considered as a prototypical giant low surface brightness (LSB) galaxy.  Malin\\,1 is among the most gas-rich galaxies known to date: M(HI)$\\approx 7 \\times 10^{10}$\\,M$_{\\odot}$ (Pickering et al. 1997) or $\\approx 5 \\times 10^{10}$\\,M$_{\\odot}$ (Matthews, van Driel \\& Monnier-Ragaigne 2001). The HI disc of the galaxy demonstrates strong non-circular motions and is strongly warped (Pickering et al. 1997; Lelli, Fraternali \\& Sancisi 2010).  Recent analysis of a Hubble Space Telescope $I$-band image suggests that Malin\\,1 has a normal barred inner spiral disc embedded in a huge diffuse LSB envelope (Barth 2007). Based on previously published photometric characteristics and HI kinematics, Seigar (2008) derived a possible mass profile for Malin\\,1. He concluded that the galaxy is baryon dominated in the centre (out to $\\sim$10 kpc) and, probably, has parameters typical of normal galaxies (of course, excluding giant low-surface brightness envelope). Recently Lelli et al. (2010) (see also Sancisi \\& Fraternali 2007) presented the similar results using their re-analysis existing HI data. The models by Seigar (2008) and Lelli et al. (2010) are based on published HI observation by Pickering et al. (1997). But these HI data are strongly affected by beam smearing and, as a result, no reliable kinematics inside central 15 kpc are currently available.  In order to investigate stellar kinematics of Malin\\,1, we have performed spectroscopic study of the galaxy with the Russian 6-m telescope. Detailed discussion and results of modeling will be published later. In this Letter, we present the most surprising result of our observations -- the discovery of the galaxy (Malin\\,1B) that probably interacts with Malin\\,1.  \\section[]{Spectral observations}   The spectroscopic data were obtained with the 6-m telescope  of the Special Astrophysical Observatory of Russian Academy of Sciences (SAO RAS) on March 30/31 2009. We observed Malin\\,1 with the multi-mode focal reducer  SCORPIO (Afanasiev \\& Moiseev 2005) in the long-slit mode with the spectral resolution about $2.5$\\AA\\, in the wavelength range 5500-6550\\AA\\, that corresponds to 5080-6050\\AA\\, in the galaxy rest frame. The studied spectral range contains numerous  absorption features (MgI, Fe, etc.) produced by old stellar population. The possible contribution of   weak  emission lines was insignificant. The slit, which length is about 6 arcmin, has been placed on the galaxy nucleus at P.A.=55$^\\circ$ (see Fig.~1). The slit width was $1\\arcsec$ with the spatial sampling $0.35\\arcsec/px$. We took in total 7200 sec exposure under atmospheric seeing $2\\arcsec$.  For data reduction and analysis we used programmes and algorithms briefly described in Zasov et al. (2008). The line-of-sight velocity and stellar velocity dispersion profiles for the stellar components have been estimated by cross-correlating galactic spectra binned along the slit with a template star HD~10380 (K3~III) spectra from the library MILES (Sanchez-Blazquez et al. 2006). We applied 2-3 pixels  binning along the slit to provide a sufficient signal-to-noise ratio, the resulting sampling was 0\\farcs70 for the inner part of the galaxy and 1\\farcs05 for several external points.  The results are presented in Fig.~2. At $r\\approx9\\arcsec$ the slit crosses the nucleus of small companion galaxy (Malin\\,1B). Fig.~2 shows that this region is kinematically decoupled from the Malin\\,1's rotation curve: it has an invert radial velocity gradient and a distinct peak at the velocity dispersion profile.  \\begin{figure} \\centering \\includegraphics[width=8.8cm, clip=]{fig1.eps} \\caption{$I$-band contour map of the central part of Malin~1. The faintest contour is about 23\\fm7/$\\square\\arcsec$, the isophotes step -- 0\\fm75/$\\square\\arcsec$. The red lines mark position of the  long-slit (1 arcsec in width). The F814W WFPC2 image was extracted from the Hubble Legacy Archive. } \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=8.8cm, clip=]{fig2.eps} \\caption{The line-of-sight velocity (top) and velocity dispersion (bottom) profiles of stars at P.A.=55$^\\circ$. Dotted lines mark photometric positions of Malin\\,1 (main galaxy) and Malin\\,1B nuclei and their systemic velocities.} \\end{figure}   \\section[]{Results and discussion}  \\subsection[]{Characteristics of Malin\\,1}  From our spectrum we determined the following parameters of the main galaxy (Malin\\,1): the heliocentric systemic velocity V$_\\mathrm{sys}$=24775$\\pm$10 km/s and the central velocity dispersion $\\sigma_0$=192$\\pm$13 km/s. Both values are in good agreement with previously published results (Table~1).  The companion galaxy -- Malin\\,1B -- has coordinates $\\alpha$(2000)=12$^h$36$^m$58.$^s$89 and $\\delta$(2000)=+14$^\\circ$19\\arcmin43\\farcs9 and is located about 9$\\arcsec$ SW from the main galaxy nucleus (Fig.~1). The systemic radial velocity of the companion is V$_{sys}$=24840$\\pm$12 km/s, the velocity difference with Malin\\,1 is 65$\\pm$16 km/s, the projected distance is 14 kpc.  Outer elliptical isophotes of Malin\\,1B are disturbed and elongated to the North-North-East (Fig.~3). The mean axial ratio the isophotes is $b/a = 0.76 \\pm 0.07$, the position angle of the major axis is P.A.=79$^\\circ$ $\\pm$ 3$^\\circ$, and the major axis is about 3$\\arcsec$ or 4.5 kpc.   \\begin{figure} \\centering \\includegraphics[width=10.5cm, angle=-90, clip=]{fig3.ps} \\caption{Top: contour map of Malin~1B from the HST $I$-band image. Orientation of the figure is the same as in Fig.~1. Isophotes are separated by factor 1.5, faintest contour is 22\\fm3/$\\square\\arcsec$. Axes are in arcsec. Bottom: photometric profile for Malin\\,1B along the apparent major axis (dots). The solid line represents approximation of the profile by standard two-component model (de Vaucouleurs' bulge and exponential disk).} \\end{figure}  The surface brightness distribution of Malin\\,1B looks usual for early-type spiral galaxies (Fig.~3). In the framework of two-component model the bulge-to-disk ratio is $\\approx1$. This value is typical for the S0-Sa galaxies (e.g. Simien \\& de Vaucouleurs 1986).  The observed luminosity of Malin\\,1B measured within $\\mu(I)=23^m/\\Box\\arcsec$ is about 1/10 of the Malin\\,1 luminosity within the same isophote. The total luminosity of the companion $M_I \\approx -20$ or $M_V \\approx -18.7$ (assuming $V-I=1.3$ for an S0 galaxy).  The observed central velocity dispersion of Malin\\,1B is $\\sigma_0$=97$\\pm$20 km/s and the galaxy parameters satisfy the Faber--Jackson relation.  The ratio $V_{max}/\\sigma_0 > 0.7$ is typical for the rotationally-supported disc galaxies, as it follows, for instance, from Binney's diagnostic $(v/\\sigma-\\epsilon)$ diagram (e.g. Kormendy 1993).  \\begin{table*} \\caption{General characteristics of the galaxies} \\begin{tabular}{lllll} \\hline \\hline Parameter &  Malin\\,1 & Malin\\,1B  & SDSS\\,J123708.91+142253.2 & Ref. \\\\ \\hline $\\alpha(2000)$     & 12$^h$36$^m$59.$^s$36 & 12$^h$36$^m$58.$^s$89 &12$^h$37$^m$08.$^s$92 \\\\ $\\delta(2000)$     & 14$^\\circ$19\\arcmin49\\farcs4 & 14$^\\circ$19\\arcmin43\\farcs9 & 14$^\\circ$22\\arcmin53\\farcs3 \\\\ &                  &    &  & \\\\ Projected separation from Malin\\,1    & 0   & 9$''$ (14 kpc) & 3\\farcm855 (350 kpc) \\\\ &                  &    &  & \\\\ Apparent magnitude & 17.6 (g) &       & 18.6 (g) & SDSS \\\\ &                  &    &  & \\\\ Heliocentric systemic velocity$^*$ (km~s$^{-1}$)& 24775$\\pm$10 & 24840$\\pm$12 & & present work \\\\ & 24800$\\pm$17 &      & 24907$\\pm$27  & SDSS \\\\ & 24819$\\pm$65 &      &               & SDSS \\\\ & 24767$\\pm$4  &  &  & Lelli et al. (2010) \\\\ & 24784$\\pm$15 &  &  & Matthews et al. (2001) \\\\ & 24755$\\pm$10 &  &  & Pickering et al. (1997) \\\\ & 24705, 24745 &  & & Impey \\& Bothun (1989) \\\\ & 24750$\\pm$10 &  & & Bothun et al. (1987) \\\\ &                  &    &  & \\\\ Central velocity dispersion (km~s$^{-1}$) & 192$\\pm$13 & 65$\\pm$16 & & present work\\\\ & 196$\\pm$15 &   &  & Barth (2007) \\\\ \\hline \\\\ \\end{tabular}  $^*$ -- conventional radial velocity obtained as $cz$ \\end{table*}   \\subsection{Malin\\,1 as an interacting galaxy}  Fig.~4 shows the large-scale spatial environment of Malin\\,1. As one can see, the galaxy is located in a relatively low-density region, near the end of an elongated structure, that is probably a filament of the large scale structure (LSS). Therefore, Malin\\,1 is within the environment typical for LSB galaxies. For example, Rosenbaum \\& Bomans (2004) and Rosenbaum et al. (2009) concluded that the LSB galaxies appear to favour the edges of filaments. This conclusion supports the idea that LSB galaxies were formed in the voids of the LSS without many galaxy interactions. Only in this case the low-density primordial fluctuations could survive. Then the galaxy would have migrated to the edges of the filaments due to gravitational infall.  \\begin{figure} \\centering \\includegraphics[width=8.cm, angle=-90, clip=]{fig4.eps} \\caption{Distribution of galaxies within a right ascension range of $160^{\\circ} \\leq \\alpha(2000) \\leq 220^{\\circ}$ (from top to bottom) and a redshift of $z\\leq0.15$ in a polar plot according to the NED database. Malin\\,1 position is indicated by red arrow. The declination range of $13.1^{\\circ} \\leq \\delta(2000) \\leq 15.6^{\\circ}$ is projected onto the plane.} \\end{figure}  There is another point of view about the origins of LSB galaxies. Some authors suggested a possible link between giant low-surface brightness galaxies and ring galaxies formed as a result of a head-on collision with a massive flyby intruder (e.g., Mapelli et al. 2008). Smoothed shell-like structures at large radii can be also produced by mergers between a disc-like galaxy and elliptical one (Pierani et al. 2010).  The collisional scenario is very attractive because Malin\\,1 has a normal inner stellar disc and a normal bulge (Barth 2007). Moreover, its extended stellar and gaseous discs can be considered as collisional rings at the late stage of their evolution. Mapelli et al. (2008) presented a model that matches in general the global photometry of Malin\\,1. Unfortunately, such a model needs the intruder with mass that is comparable with the total mass of the main galaxy or only a few times less. Malin\\,1B looks too small and unsufficiently massive to produce an extended ring-like structure. It could have been more massive originally but has been tidally stripped during the interaction with Malin\\,1. But to our mind, its position and relative velocity favour a close orbit rather than a flyby at escape speed that is commonly preassumed when one considers a collisional origin of ring-like galaxies (Lynds \\& Toomre 1976; Hernquist \\& Weil 1993; Horellou \\& Combes 2001).  Nevertheless Malin\\,1B could cause the appearance of a single dusty spiral arm in Malin\\,1 that has been recently revealed (Moore \\& Parker 2006). This type of structure is often considered as the result of an interaction with a satellite on a retrogade orbit (Athanassoula 1978).  Malin\\,1B could also produce two-armed structure and trigger a bar instability in the inner disc of Malin\\,1. The inner spiral structure is definitely seen in the Hubble Space Telescope $I$-band image (Figs.~1,2 of Barth 2007) and the bar can be clearly recognized by analysing the surface brightness distribution of Malin\\,1 (Fig.~3 of Barth 2007). But we definitely need a more massive companion to support the scenario by Mapelli et al. (2008)  The nearest bright galaxy with known redshift is SDSS\\,J123708.91+142253.2 with $V_\\mathrm{hel} = 24907 \\pm 27$ (Table~1). The galaxy is located at the projected distance of 3.9 arcmin or 350 kpc far from Malin\\,1 and has an apparent magnitude $g(\\mathrm{SDSS})=18.6$. The velocity difference between this galaxy and Malin\\,1 is rather small --- 132$\\pm$29 km/s.  We can use parameters of the collisional model presented by Mapelli et al. (2008) to check the galaxy SDSS\\,J123708.91+142253.2 as a candidate for being a possible intruder. The ratio of the Malin\\,1 and  ``intruder'' optical luminosities (that is about $1/2.5$ in the $g$-band) is in agreement with the value of the mass ratio of target and intruder galaxies in the model under discussion.  The collisional model (Mapelli et al. 2008) suggests an impact at an inclination angle of $\\sim 14^\\circ$ with respect to the angular momentum axis of the target (Malin\\,1 in our case). According to Moore \\& Parker (2006) the inclination angle of a stellar disc of Malin\\,1 is $\\sim 45^\\circ$. Then the direction of intruder motion would have an inclination angle to the line of sight between $\\theta \\sim 30^\\circ$ and $\\theta \\sim 60^\\circ$. Let us choose the largest value of the angle for further estimates. The initial position and the velocity (with respect to the centre-of-mass of the progenitor of the ring galaxy) for the intruder adopted by Mapelli et al. (2008) are $r_\\mathrm{} \\approx 33$ kpc, $v_\\mathrm{ini} \\approx 900$ km/s. It implies that the intruder radial velocity would be of $v = v_\\mathrm{ini} \\cos{\\theta} \\sqrt{r_\\mathrm{ini} / r} \\sim 130$~km/s at the projected distance $r_\\mathrm{proj} = r \\sin{\\theta} = 350$ kpc. Both estimates mean that it takes the intruder about 1~Gyr ($t \\approx \\frac{2}{3} r / v$) to get there, providing a very small impact parameter and a flyby as suggested by Mapelli et al. (2008).  Our simple velocity estimate for the intruder is in agreement with the data presented here. Moreover, current location of the galaxy SDSS\\,J123708.91+142253.2 maintains the idea for the closest passage to happen 1~Gyr ago that is also in accordance with the model by Mapelli et al. (2008). But this otherwise attractive scenario has one serious drawback.  The model by Mapelli et al. (2008) predicts the existence of strong non-circular motions that could be associated with the expansion velocity of the old ring and/or with the fallback of a part of the ejected matter towards the centre. The strength of non-circular motions can be quantified by the ratio of radial and tangential velocities of gas. Simulations show that there is a region in the outer part of the galaxy ($\\sim$70--80 kpc far from the center) where radial velocities are comparable with tangential ones (Fig.~12 of Mapelli et al. 2008). However, the analysis of HI data (to be discussed in a forthcoming paper) shows that possible radial motions are smaller (see also Lelli et al. 2010). Namely, the radial component of non-circular gas motions in the galaxy outskirts have a maximal values about $0.3-0.4$ of rotation velocities. Moreover, the azimuthal distribution of non-circular motions is more complex than simple inflow/expansion pattern predicted by numerical simulations.  In any case, the structure of Malin\\,1 seems  more complex than predicted from Mapelli et al. (2008) model because the real situation should include two events: a possible strong head-on collision with a massive companion (SDSS\\,J123708.91+142253.2) and a current interaction with smaller Malin\\,1B galaxy.  Thus, we do not reject the collisional model of Malin\\,1 but it definitely needs to be refined, as far as, new observational data are urgently needed. ",
        "conclusions": "In this letter we present the results of spectroscopic observations of the galaxy Malin\\,1 and its newly discovered companion Malin\\,1B. Our main conclusions are  \\begin{enumerate}  \\item Malin\\,1 is currently undergoing a minor merger with a relatively small companion galaxy. This galaxy is at the projected distance of 14 kpc and shows a small velocity difference with Malin\\,1 systemic velocity ($\\Delta$V=65 km/s).  \\item Accreting companion might be responsible for the main morphological features in the central part of Malin\\,1. It could produce a two-armed inner structure and trigger the bar instability (Fig.~1). Also, one can relate this companion with one-armed spiral pattern in Malin\\,1 disc (see Fig.~4 in Moore \\& Parker 2006).  \\item Malin\\,1 is in a relatively low-density large scale spatial environment that is typical for LSB galaxies.  \\item We discuss the possible origins of Malin\\,1 global structure due to a bygone head-on collision with a massive intruder (mechanism proposed by Mapelli et al. 2008). Available data do not contradict this scenario (for instance, we even identified possible intruder galaxy SDSS\\,J123708.91+142253.2) but more detailed simulations and new observational data are needed for definite conclusions.  \\end{enumerate}"
    },
    "1005/1005.2157_arXiv.txt": {
        "abstract": "{} {We wish to separate and quantify the CO luminosity and CO-\\HH\\ conversion factor applicable to diffuse but partially-molecular ISM when \\HH\\ and CO are present but C\\p\\ is the dominant form of gas-phase carbon.} {We discuss galactic lines of sight observed in \\HI, \\hcop\\ and CO where CO emission is present but the intervening clouds are diffuse (locally \\AV\\ $\\la 1$ mag) with relatively small CO column densities $\\NCO \\la 2\\times10^{16}\\pcc$.  We separate the atomic and molecular fractions statistically using \\EBV\\ as a gauge of the total gas column density and compare \\NHH{} to the observed CO brightness.} {Although there are \\HH-bearing regions where CO emission is too faint to be detected, the mean ratio of integrated CO brightness to \\NHH{} for diffuse ISM does not differ from the usual value of 1\\Kkms\\ of integrated CO brightness per $2\\times10^{20}$ \\HH $\\pcc$ .  Moreover, the luminosity of diffuse CO viewed perpendicular to the galactic plane is 2/3 that seen at the Solar galactic radius in surveys of CO emission near the galactic plane.} {Commonality of the CO-\\HH\\ conversion factors in diffuse and dark clouds can be understood from considerations of radiative transfer and CO chemistry.  There is unavoidable confusion between CO emission from diffuse and dark gas and misattribution of CO emission from diffuse to dark or giant molecular clouds.  The character of the ISM is different from what has been believed if CO and \\HH\\ that have been attributed to molecular clouds on the verge of star formation are actually in more tenuous, gravitationally-unbound diffuse gas.} ",
        "introduction": "It is a truism that sky maps of CO emission are understood as uniquely tracing the Galaxy's molecular clouds, dense, cold strongly-shielded regions where the hydrogen is predominantly molecular and the dominant form of gas phase carbon is CO. Moreover, CO emission plays an especially important role in ISM studies as the tracer of cold molecular hydrogen through the use of the so-called CO-\\HH\\ conversion factor which directly scales the integrated \\cotw\\ J=1-0 brightness \\WCO{} to \\HH\\ column density \\NHH{}.  This deceptively simple conversion is critically important to determining molecular and total gas column densities and so to understanding the most basic properties of star formation \\citep{LerWal+08,BigLer+08,BotKen+09}, the origins of galactic dust emission \\citep{DraDal+07}, and other such fundamentals.  Yet, it is increasingly recognized that CO emission is present along lines of sight lacking high extinction or large molecular column densities \\citep{LisLuc98}.  It is also the case that some very opaque lines of sight showing CO emission are comprised entirely of diffuse material and \\HH-bearing diffuse clouds \\citep{McCHin+02}:  a discussion of such a line of sight from our own work is described in Appendix A here.  Even in canonical dark clouds like Taurus, detailed high-resolution mapping of the CO emission \\citep{GolHey+08} reveals that much of it originates in relatively weakly-shielded gas where \\coth\\ is strongly enhanced through isotopic fractionation, implying that the dominant form of gas phase carbon must be C\\p\\ \\citep{WatAni+76}.  \\begin{figure*} \\psfig{figure=14510f1.eps,height=8cm} \\caption[]{Atomic and molecular absorption and emission vs. total reddening.  Left: Integrated VLA \\HI{} optical depth from \\cite{DicKul+83} and this work.  Middle: integrated PdBI \\hcop\\ optical depth from \\cite{LucLis96} and this work.  Right: integrated ARO12m CO J=1-0 brightness at 1\\arcmin\\ resolution.  In each case the horizontal axis is the total line of sight reddening \\EBV\\ \\citep{SchFin+98}. For explanation of the symbols used in the plots, see Sect. 3.} \\end{figure*}  Conversely, it is also the case that molecular gas is detected in the local ISM even when CO emission is not.  Lines of sight with $\\NCO \\ga 10^{12} \\pcc$, $\\NHH \\ga 10^{19} \\pcc$ have long been detectable in surveys of uv absorption \\citep{SonWel+07,BurFra+07,SheRog+07,SheRog+08}, with expected integrated CO brightnesses as low as $\\WCO = 0.001\\Kkms$ \\citep{Lis07}.  And, as discussed here, mm-wave \\hcop\\ and CO absorption from clouds with $\\NHH \\ga 10^{20}~\\pcc$ are also more common than CO emission along the same lines of sight \\citep[see][and Appendix A]{LucLis96,LisLuc00}.  Thus we are led to ask two questions that are of particular importance to the use of CO emission as a molecular gas tracer.  First, where and how does the observed local CO luminosity really originate? Second, how completely is the molecular material represented by CO emission?  The approach we take to address these issues is based on radiofrequency surveys of \\HI{}, \\hcop\\ and CO absorption and emission along lines of sight through the Galaxy toward extragalactic background sources.  By combining 1) measurements of extinction (constraining the total gas column density), 2) measurements of \\HI{} absorption (to determine the gas column of atomic hydrogen), 3) the strength of \\hcop\\ absorption (tracing \\HH\\ directly) and 4) the integrated CO J=1-0 brightness \\WCO, we relate \\WCO\\ to \\NHH{} along sightlines where we have previously shown that the intervening gas is diffuse, neither dark nor dense, and the CO column densities are relatively small.  The results are somewhat surprising: although there is much variability, the mean CO brightness per \\HH-molecule \\WCO/\\NHH{}, i.e. the CO-\\HH\\ conversion factor, does not differ between diffuse and fully molecular clouds.  Although this was phenomenologically inferred long ago, the physical basis for it is now better understood in terms of the radiative transfer and chemistry of \\HH- and CO-bearing diffuse and dark gas.  The plan of the present work is as follows.  Section 2 describes the observational material that is used here, some of which is new.  Section 3 derives the CO-\\HH{} conversion factor in diffuse gas.  Section 4 discusses the fraction of the local galactic CO luminosity (viewed perpendicular to the galactic plane) that can be attributed to diffuse gas. Section 5 discusses the physical processes at play to set the ratio of CO brightness to \\HH\\ column density and explains why the same value may apply to dark and diffuse gas. Section 6 discusses which molecular emission diagnostics might actually be used to distinguish between the CO contributions from diffuse and dark gas. Sections 7 and 8 present a brief summary and discuss how our concept of the ISM might change when a substantial portion of the observed CO emission is ascribed to diffuse rather than dense molecular gas. ",
        "conclusions": "In Sect. 2 and 3 we described and considered a sample of lines of sight studied in \\HI{} and molecular absorption and known to be comprised of diffuse gas.  Their molecular component shows features whose CO, \\hcop\\ and other molecular column densities are small compared to those of dark clouds (in the case of CO, at least 30 times smaller).   There is often quite substantial fractionation of \\coth\\ (indicating that the dominant form of carbon is C\\p) and the rotational excitation of CO is sub-thermal with implied cloud  temperatures typical of those determined directly for diffuse \\HH\\ in optical/uv surveys, i.e. 30 K or more.  Using an externally-determined value for the ratio of total \\HI{} column density to integrated \\HI{} absorption and the standard equivalence between reddening and \\NH{} we derived the molecular gas fraction for this sample to be \\fHH\\ = 0.35, in the middle of the range of other estimates for the diffuse ISM as a whole based on optical (mainly CH) and uv (\\HI{} and \\HH) absorption studies.  We showed that this estimate for \\fHH\\ implies the same value $\\XHCOp = 3\\times 10^{-9}$ that was previously determined from comparisons of OH and \\hcop\\ column densities in individual clouds.  We then compared measured CO brightnesses with the inferred molecular gas column densities to derive the ensemble mean $\\NHH/\\WCO$ conversion factor.  Surprisingly, We found this mean to be just equal to the locally-accepted value $2.0\\times 10^{20} \\HH$ /(\\Kkms) for ``molecular'' gas believed to reside in dense dark fully-molecular clouds near the galactic equator.  Such exact agreement is probably something of an accident of sampling, but the fact that diffuse and dark gas would have very similar $\\NHH/\\WCO$ conversion factors, which had been inferred empirically long ago, now has a firmer physical basis.  In Sect. 5 we explained it as the result of the brightening of CO J=1-0 emission per CO molecule that was theoretically predicted for warmer more diffuse gas by \\citet{GolKwa74}, which compensates for the lower relative abundance \\XCO{} there.  The mean CO abundance observed in optical absorption in diffuse clouds $\\mean{\\XCO} = 3\\times 10^{-6}$, combined with the observed and expected brightness per CO molecule, \\WCO/\\NCO{} = 1 \\Kkms/$10^{15} {\\rm CO}\\pcc$, can be be combined to form an CO-\\HH\\ conversion factor of \\NHH{}/\\WCO $= 10^{15}/3\\times 10^{-6} = 3.3 \\times 10^{20}\\pccpKkms{\\HH}$.  In Sect. 4 we derived the expected brightness of diffuse gas viewed perpendicular to the galactic plane from afar, 0.47\\Kkms, and compared that to the value expected from surveys of CO emission in the galactic plane, combined with a narrow (60 pc dispersion) Gaussian vertical distribution; that is 0.75\\Kkms.  This suggests that there has been confusion in the general attribution of CO emission to ``molecular clouds'' when in fact much of it arises in the diffuse ISM.  This view is consistent with the motivations discussed in the Introduction, whereby CO emission is increasingly being found along lines of sight lacking high extinction and whereby CO emission seen along dark lines of sight are found (through molecular absorption studies and in other ways) to originate in components having the relatively small molecular number and column densities typical of diffuse gas.  An example of such a line of sight is given in Appendix A here.  We briefly discussed in Sect. 4 the decomposition of the total mean density of neutral gas in the nearby ISM, 1.2 H $\\pccc$ \\citep{Spi78}, into its atomic and molecular constituents.  We noted that although the balance is generally believed to be roughly 50-50 \\citep{Cox05}, some emission might shift to the diffuse side of the balance sheet if CO emission is reinterpreted.  Moreover, we pointed out that the molecular contribution to the true local mean density from large-scale galactic CO surveys in the galactic plane should be questioned more generally because it is unclear to what extent Spitzer's estimate, based on the earlier optical work of M\\\"unch, incorporates the contribution of optically-opaque gas.  Although the ability to discriminate between the separate contributions to \\WCO\\ from diffuse and darker, denser gas is limited when only \\cotw\\ is considered, it should be possible to infer the nature of the host gas using other emission diagnostics (see Sect. 6).  The most efficient of these is probably the brightness of \\coth, which, although enhanced by fractionation, is still substantially weaker, relative to \\WCO, in diffuse gas.  Searching for emission from species having higher dipole moments such as CS J=2-1 and HCN (and probably not \\hcop\\ because it is chemically so ubiquitous and more easily excited) are alternatives that might require somewhat longer integration times."
    },
    "1005/1005.0014_arXiv.txt": {
        "abstract": "The Fundamental Plane of early type galaxies is observed to have finite thickness and to be tilted from the virial relation.  Both of these represent departures from the simple assumption that dynamical mass-to-light ratios ($M_{dyn}/L$) are constant for all early type galaxies.  We use a sample of 16,000 quiescent galaxies from the Sloan Digital Sky Survey to map out the variations in $M_{dyn}/L$ throughout the 3D Fundamental Plane space defined by velocity dispersion ($\\sigma$), effective radius ($R_e$), and effective surface brightness ($I_e$).  Dividing $M_{dyn}/L$ into multiple components allows us to separately consider the contribution to the observed $M_{dyn}/L$ variation due to stellar population effects, IMF variations, and variations in the dark matter fraction within one $R_e$.  Along the FP, we find that the stellar population contribution given some constant IMF ($M_{\\star,IMF}/L$) scales with $\\sigma$ such that $M_{\\star,IMF}/L \\propto f(\\sigma)$.  Meanwhile, the dark matter and/or IMF contribution ($M_{dyn}/M_{\\star,IMF}$) scales with $M_{dyn}$ such that $M_{dyn}/M_{\\star,IMF} \\propto g(M_{dyn})$.  This means that the two contributions to the tilt of the FP rotate the plane around different axes in the 3D space.  The observed tilt of the FP requires contributions from both, with dark matter/IMF variations likely comprising the dominant contribution.  Looking at $M_{dyn}/L$ variations through the {\\it thickness} of the FP, we find that $M_{dyn}/L$ variations must be dominated either by IMF variations or by real differences in dark matter fraction with $R_e$.  This means that the finite thickness of the FP is due to variations in the stellar mass surface density within $R_e$ ($\\Sigma_{\\star,IMF}$), not the fading of passive stellar populations and it therefore represents genuine structural differences between early type galaxies.  These structural variations are correlated with galaxy star formation histories such that galaxies with higher $M_{dyn}/M_{\\star,IMF}$ have higher [Mg/Fe], lower metallicities, and older mean stellar ages.  We discuss several physical mechanisms that might explain the observed co-variation between $M_{dyn}/M_{\\star,IMF}$ and galaxy star formation histories.  It is difficult to explain the observed enhancement of $\\alpha$-elements in lower-surface-brightness galaxies by allowing the IMF to vary.  Differences in dark matter fraction can be produced by variations in the ``conversion efficiency'' of baryons into stars or by the redistribution of stars and dark matter through dissipational merging.  The former explanation, specifically a model in which some galaxies experience low conversion efficiencies due to premature truncation of star formation, provides a more natural explanation for the co-variation of $M_{dyn}/M_{\\star,IMF}$ and the observed stellar population properties. ",
        "introduction": "Early type galaxies are observed to obey many scaling relations among their structural properties.  Early work identified a number of 1D relations, such as the Faber-Jackson relation between galaxy luminosity ($L$) and velocity dispersion ($\\sigma$) \\citep{faber76}, variations in galaxy mass-to-light ratio ($M/L$) versus $L$ (\\citealt{tinsley81}; \\citealt{faber87}), and various correlations of $\\sigma$ and $L$ with galaxy effective radius ($R_e$) and effective surface brightness ($\\mu_e$) \\citep{kormendy85}, or with galaxy core radius ($r_c$) and central surface brightness ($I_{0}$) \\citep{lauer85}.  These relations generally reflect a 1D mass sequence of galaxies.  It is clear however that early type galaxies comprise at least a two-parameter family in terms of their structure.  In the three-dimensional parameter space of galaxy properties defined by $\\sigma$, $R_e$, and surface brightness (expressed as $\\mu_e$ in magnitudes or as $I_e$ in linear flux units), early type galaxies populate a relatively tight two-dimensional plane, known as the Fundamental Plane (FP, \\citealt{djorgovski87, dressler87}). Projections of the FP appear narrow in some orientations, leading to the seemingly 1D relations listed above.  The FP can be understood as a manifestation of the virial plane predicted for relaxed systems, under the assumption that galaxy mass-to-light ratios ($M_{dyn}/L$) are constant (or at least smoothly varying) for all galaxies.  If $M_{dyn}/L$ were strictly constant and if structures were homologous for all early type galaxies, the FP would be equivalent to the virial plane, which takes the form $R_e \\propto \\sigma^2 I_e^{-1}$, and would be infinitely thin.  Instead, the FP is rotated from the virial plane (e.g., \\citealt{dressler87, djorgovski87}).  \\citet{jorgensen96} find $R_e \\propto \\sigma^{1.24} I_e^{-0.82}$ for the FP of local cluster early type galaxies, as measured in the Gunn $r$ band.  In addition, the FP has finite thickness (e.g., \\citealt{jorgensen96, forbes98, wuyts04, hyde09b, gargiulo09}) indicating further departures from a single virial plane.  The ``tilt'' of the FP and its finite thickness indicate departures from the simplistic assumption that $M_{dyn}/L$ is constant for all early type galaxies and that their structures are homologous.  This is the third in a series of four papers that explores the mapping between the 2D family of early type galaxy stellar populations and their structural properties.  The first three papers in the series systematically gather together a number of galaxy properties that vary through the different dimensions of galaxy parameter space. \\citet[hereafter Paper I]{graves09_paperI} showed that stellar populations vary differently as functions of $\\sigma$ than as functions of $L$.  This systematic variation suggests an underlying multi-dimensional parameter space.  \\citet[hereafter Paper II]{graves09_paperII} mapped out the second dimension of stellar population properties in 3D Fundamental Plane space and correlated stellar population variations with FP parameters.  In this paper, we focus on the corresponding mass-to-light ratios and how they relate to the distribution of galaxies in FP space.  We show here that the stellar population variations measured in Paper II do not provide enough variation in galaxy mass-to-light ratios to explain either the tilt or the thickness of the FP.  Some correlated variation is also required in either the initial mass function (IMF) with which galaxies form their stars, or in the dark matter fraction inside $R_e$.  At the same time, these variations are a key component of structural differentiation in galaxies that must also be related to their star formation histories.  The results of these first three papers are gathered together in Paper IV (Graves et al. 2010, submitted to ApJ), which uses them to propose scenarios for galaxy formation and evolution through the thickness of the FP.  Section \\ref{data} briefly outlines the sample of galaxies used in this work, as well as our methods for grouping together similar galaxies to produce high signal-to-noise spectra that span the space of galaxy properties.  In section \\ref{ml_contributions}, we discuss the various physical mechanisms that may contribute to the observed variations in $M_{dyn}/L$ throughout the FP.  Some of this variation is due to stellar population effects, but quantifying these depends critically on the stellar population models used to measure stellar mass-to-light ratios ($M_{\\star,IMF}/L$).  We therefore investigate multiple ways of measuring $M_{\\star,IMF}/L$ from stellar population models in section \\ref{mstar_l}.  This comparison allows us to explore the range of possible $M_{\\star,IMF}/L$ values for our sample galaxies and to understand the biases inherent in each model.  We demonstrate that, regardless of which method is chosen to measure $M_{\\star,IMF}/L$, the observed stellar population variations are inadequate to explain either the tilt or the thickness of the FP. Section \\ref{ml_in_3d} shows that the two contributions to the FP tilt (1: from known stellar population variations and 2: from IMF and/or dark matter fraction variations within $R_e$) rotate the plane around different axes in 3D Fundamental Plane space.  It also investigates the various contributions to the {\\it thickness} of the FP and shows that the spread in $M_{dyn}/L$ through the plane is due primarily to the second factor: variations in the IMF and/or dark matter fractions. Section \\ref{sfh} discusses physical models that may produce the observed correlations between dark matter fraction or IMF variation and galaxy star formation histories.  It illustrates that the observed trends can be explained by a model in which some galaxies experience low ``conversion efficiency'' turning baryons into stars due to trucated star formation histories.  Finally, section \\ref{conclusions} summarizes our conclusions.  \\begin{figure*}[t] \\includegraphics[width=1.0\\linewidth]{f1_small.eps} \\caption{Sorting and binning galaxies in 3D FP space.  (a)~An edge-on view of the best fit to the FP, using a least-squares fit of $I_e$ onto $R_e$ and $\\sigma$, which gives $f(\\sigma,R_e) = 1.16 \\log \\sigma - 1.21 \\log R_e + 0.55$.  We divide the FP into three slices based on surface brightness residuals ($\\Delta \\log I_e$) to define low-surface-brightness, midplane, and high-surface-brightness slices. (b)~Within each slice, galaxies are further divided into 6 bins in $\\log \\sigma$ and 5 bins in $\\log R_e$, as illustrated. Diagonal shading indicates the six bins that do not contain enough galaxies to produce robust mean spectra and are therefore excluded from our analysis.  Vertical shading indicates a bin where stellar population ages cannot be reliably measured (see Paper II). }\\label{fp_bins} \\end{figure*} ",
        "conclusions": "In Paper II, we demonstrated that the stellar population properties, and therefore the star formation histories, of quiescent early type galaxies form a two-parameter family.  These map onto a cross-section through the Fundamental Plane.  In this paper, we have explored the associated mass-to-light variations in 3D FP space.  This has allowed us to map out the various components of mass-to-light variation both along the FP and through its thickness, with the following results:  \\renewcommand{\\labelenumii}{\\alph{enumii})} \\begin{enumerate} \\item{We map $M_{dyn}/L$, $M_{dyn}/M_{\\star,IMF}$, and $M_{\\star,IMF}/L$ throughout 3D Fundamental Plane space.  These quantities can be well-approximated by first-order functions (i.e., hyperplane fits) of $\\log \\sigma$, $\\log R_e$, and $\\Delta \\log I_e$.  These fits are provided in Table \\ref{fits_table}.} \\item{We confirm that the variation in $M_{\\star,IMF}/L$ due to stellar population variations cannot account for the tilt of the FP, as demonstrated by many previous authors.  Variations in the dark matter fraction within $R_e$ (i.e., $M_{dyn}/M_{\\star}$) or in the IMF (i.e., $M_{\\star}/M_{\\star,IMF}$) are required to explain the tilt of the FP.  Using the $M_{\\star,IMF}$ estimates of \\citet{gallazzi05}, stellar population effects contribute only 1/3 of the tilt of the FP, while dark matter/IMF effects contribute the remaining 2/3.} \\item{The stellar population and the dark matter/IMF contributions to the FP tilt rotate the plane around different axes in the 3D space, with $M_{\\star,IMF}/L \\propto f(\\sigma)$, while $M_{dyn}/M_{\\star,IMF} \\propto g(M_{dyn})$. } \\item{Although Paper II showed that stellar populations vary through the thickness of the FP (in agreement with \\citealt{forbes98}, \\citealt{wuyts04}, and \\citealt{gargiulo09}), the associated variations in $M_{\\star,IMF}/L$ alone {\\it cannot account for the thickness of the FP}.  Variations in the dark matter contribution within $R_e$ (i.e., $M_{dyn}/M_{\\star}$) or in the IMF (i.e., $M_{\\star}/M_{\\star,IMF}$) are required to explain the thickness of the FP (as well as the tilt).  Using the $M_{\\star,IMF}$ estimates of \\citet{gallazzi05}, stellar population effects contribute only 22\\% of the thickness of the FP, while dark matter/IMF effects contribute the remaining 78\\%.} \\item{Because stellar population effects contribute only $\\sim 22$\\% of the thickness of the FP, the observed $M_{dyn}/L$ variations through the plane are dominated by variations in $\\Sigma_{\\star,IMF}$. This means that galaxies do not move below the plane by fading; they either form initially with lower $\\Sigma_{\\star,IMF}$ or move there due to {\\it structural} changes in the stellar mass distribution.  } \\item{Combining the results of this work with the results of Paper II, we find that variations in the dark matter/IMF contribution through the thickness of the FP are associated with differences in galaxy star formation histories.  The correlation is such that galaxies with higher $M_{dyn}/M_{\\star,IMF}$ also have higher [Mg/Fe], older ages, and lower metallicities than their lower-$M_{dyn}/M_{\\star,IMF}$ counterparts at the same $\\sigma$, which may indicate that the higher-$M_{dyn}/M_{\\star,IMF}$ galaxies have experienced shorter duration star formation.  Possible physical mechanisms for producing high-$M_{dyn}/M_{\\star,IMF}$ and the associated stellar population variations include: \\begin{enumerate} \\item{A top-heavy IMF in some galaxies, which produces more SNe II relative to SNe Ia and therefore higher [Mg/Fe].  However, this scenario does not simultaneously produce the observed lower [Fe/H] and [Mg/H] in these galaxies.  Invoking IMF variations to explain the observed $M_{dyn}/M_{\\star,IMF}$ variations would require an additional physical mechanism, such as enhanced SN feedback that removes a substantial fraction of the metal-enriched material.} \\item{Higher dark matter fraction within $R_e$ due to low-dissipation merging in galaxies that have already quenched their star formation. This scenario is plausible, but the predicted trend with stellar population age seems backwards to expectations since early-forming galaxies are naturally more gas-rich, not less.} \\item{Low conversion efficiency of baryons into stars through the premature truncation of star formation.  There are various possible mechanisms for this trunction, including ram-pressure stripping in satellite galaxies, powerful AGN feedback, massive halo quenching, or low-mass progenitor galaxies that are strongly affected by supernova feedback (as compared to galaxies that form most of their stars {\\it in situ} in deep potential wells).  This truncation scenario is appealing because it simulatenously matches the lower stellar masses, older ages, higher [Mg/Fe], lower [Fe/H], and lower [Mg/H] observed in the high-$M_{dyn}/M_{\\star,IMF}$ galaxies at fixed $\\sigma$.} \\end{enumerate} } \\item{There is an intriguing parallel between the known increase of $M_{halo}/M_{\\star}$ in large haloes and the increase in $M_{dyn}/M_{\\star,IMF}$ along the FP observed in this galaxy sample. This suggests that the conversion efficiency of baryons into stars inside $R_e$ mimics the overall conversion efficiency for the parent halo.  The truncation of star formation in massive galaxies is widely accepted as the explanation for the increase in $M_{halo}/M_{\\star}$ with galaxy mass, although the truncation {\\it mechanism} is debated.  It may be that this same truncation mechanism also produces differences in conversion efficiency {\\it inside} $R_e$, and that this effect drives the observed variation in $M_{dyn}/M_{\\star,IMF}$ through the thickness of the FP as well.} \\end{enumerate}  These results have focused on quiescent galaxies only, excluding early type galaxies with ongoing star formation, Seyfert, or LINER activity. Seyfert and LINER hosts are known to have systematically younger stellar population ages than their quiescent counterparts \\citep{kauffmann03_agn, graves07, schawinski07}, but further analysis is required to determine whether they follow the other stellar population and $M/L$ trends explored here for quiescent galaxies.  Taken as an ensemble, the results presented in this series of papers illustrate a connection between the star formation histories and the present-day mass structures of early type galaxies.  As such, they represent a stringent test for models of galaxy formation and a challenge for the next generation of semi-analytic models.  They also highlight the emerging importance of abundances in testing model star formation histories.  Paper IV makes extensive use of this abundance information to expand on and quantify the premature truncation model proposed here."
    },
    "1005/1005.4960_arXiv.txt": {
        "abstract": "{Recent observational results  found a bend in the Tully-Fisher Relation in such a way that low mass systems lay below the linear relation described by more massive galaxies.} {We intend to investigate the origin of the observed features in the stellar and baryonic Tully-Fisher relations and analyse the role played by galactic outflows on their determination.} {Cosmological hydrodynamical simulations which include Supernova feedback were performed in order to follow the dynamical evolution of galaxies.} {We found that Supernova feedback is a fundamental process  in order  to reproduce the observed trends in the stellar Tully-Fisher relation. Simulated slow rotating systems tend to have lower stellar masses than those predicted by the linear fit to the massive end of the relation, consistently with observations. This feature is not present  if Supernova feedback is turned off. In the case of the baryonic Tully-Fisher relation, we also detect a weaker tendency for smaller systems to lie below the linear relation described by larger ones. This behaviour arises  as a result of the more efficient action of Supernovae in the regulation of the star formation process and in the triggering of powerful galactic outflows in shallower potential wells which may heat up and/or expel  part of the gas reservoir. } {} ",
        "introduction": "The origin of the Tully-Fisher relation \\citep{tf77} has been analysed by numerous observational and theoretical works since it provides important constraints for galaxy formation models \\citep[e.g.][]{vd98, mo98, avila09}. Originally defined as the relation between the luminosity and the rotation velocity for spiral galaxies, it is now accepted that this is a proxy of a more fundamental one between stellar mass and rotation velocity \\citep[e.g.][]{cons05, flo06, cres09} or even between baryonic mass and rotation velocity \\citep[e.g.][]{bdj01, ver01, guro04, mcga05, puech10}  of the form $M \\propto V^{\\alpha}$. Observational evidence suggests the existence of a break in the stellar Tully-Fisher relation (sTFR) at a rotation velocity of $ \\sim 90 \\, {\\rm km \\, s^{-1}}$ \\citep{mcg00}. Even more, recently \\citet{mcg10} also reported a break  in the baryonic Tully-Fisher relation (bTFR) but at a lower velocity.   Assuming that galaxies formed in virialized haloes and that the fraction of gas transformed into stars is independent of  halo mass,  a relation between stellar mass and circular velocity  with a slope of $\\alpha \\sim 3$ is directly obtained \\citep{wf91, mo98}. This theoretical prediction is reproduced in cosmological simulations when non effective mechanisms to regulate the gas cooling and star formation activity  are introduced \\citep{tis97, sn99}. Therefore,  changes in the slope of this relation to match observed reported values, which  are on average close to  $\\alpha \\sim 4$, might require the action of other physical processes. Theoretical  findings suggest that the triggering of powerful mass-loaded galactic winds by massive Supernova (SN) explosions would have an efficiency which depends on the virial halo so that larger effects are expected in low mass systems \\citep[][]{larson04}. In fact, \\citet{ds86}  estimated analytically that haloes with circular velocity smaller than $\\sim 100$ km s$^{-1}$  would be strongly affected by massive SN events. Nonetheless, the role of galactic outflows on the origin of the relations between the stellar and dynamical properties of dwarf galaxies still constitutes an open problem \\citep{maclow99, tassis08}. In a  cosmological framework, only recently it has been possible to start simulating  the triggering of galactic outflows on more physical basis  \\citep[e.g.][]{ scan06, stin07}. In fact, cosmological simulations that include this process obtained  sTFRs in more qualitative agreement with observations \\citep{gov07, croft09}. There is evidence that the implementation of a suitable feedback model could account for a bend in the TFR as suggested by semi-analytical works \\citep[e.g.][]{kang05, nag05, vdbosh09, guo09} and some hydrodynamical simulations \\citep[e.g.][]{tassis08}. However, a detailed analysis of the effects of SN feedback on the TFR using a physically motivated model was still missing.   In this paper, we present results on the analysis of the sTFR and bTFR by using a SN feedback model specially suitable to study the formation of galaxies in cosmological scenarios \\citep{scan05, scan06} since it does not include any scale-dependent parameter and yet, it is capable to produce the expected anti-correlation between the strength of the galactic outflows and the potential well of the haloes,  as shown by \\citet{scan06}.   This paper is organized as follows. Section \\ref{sec:simus} describes the numerical experiments. Section \\ref{sec:results} outlines the results. In Section \\ref{sec:discussion} we discuss about our main findings and in Section \\ref{sec:conclusions} we present the conclusions of this work. ",
        "conclusions": "\\label{sec:conclusions}  We have studied the sTFR and bTFR by performing hydrodynamical simulations in a cosmological scenario including the action of the combined multiphase and SN feedback model of \\citet{scan05, scan06}. Our results suggest the existence of a change in the slope of the sTFR and bTFR produced by the action of SN feedback. At $z=0$, the simulated sTFR is very tightly defined  and  in general agreement with observations  \\citep[e.g.][]{mcg00}. We detect a systematical change in the slope of the sTFR from  approximately  $\\sim 100 \\ {\\rm km \\ s}^{-1}$, at least since $z=3$, so that lower velocity systems tend to have smaller stellar masses than those predicted by the linear regression of the fast rotators. The comparison with higher and lower  resolution runs shows that this feature is not just a numerical artefact. Regarding the local bTFR, over the velocity range resolved by our simulations, it can be fitted by a single linear model. At higher redshifts, we  detect a weak trend for a change in the slope. Similar trends are found by using $V_{\\rm max}$ as the kinematical estimator.   Our results suggest an evolution of $\\approx 0.44$ dex from $z=3$ to $z=0$ for the sTFR  which is global agreement with that reported by \\citet{cres09} of $0.41 \\pm 0.11$ dex from $z \\approx 2.2$ to $z=0$. For the baryonic relation, our model predicts a slightly weaker evolution of  $\\approx 0.30$ dex from $z=3$ to $z=0$. In the wind-free run, both the sTFR and the bTFR show the same level of evolution ($\\approx 0.55$ dex) as expected since baryons are efficiently transformed into stars as soon as they collapse within the potential well of a galaxy.  The characteristic velocity of $\\sim 100 \\ {\\rm km \\ s}^{-1}$ appears {\\it naturally } as a consequence of the physically motivated multiphase and SN feedback treatment adopted in these simulations. This model is successful at establishing a self-regulated star formation process without using any global properties of the systems but resorting to the local estimations of the thermodynamical properties of the gas  in the distinct components of the ISM. We have checked that this result is robust against variations of the SN input parameters. Nevertheless, we acknowledge the existence of several issues that need to be solved  in the future such as the level of change in the slope which seems to be weaker in our simulations compared to observations. In this regard, our findings suggest that the study of the sTFR and bTFR at different cosmic times could help to constrain SN feedback models."
    },
    "1005/1005.1068_arXiv.txt": {
        "abstract": "The study of short-duration gamma-ray bursts (GRBs) experienced a complete revolution in recent years thanks to the discovery of the first afterglows and host galaxies starting in May 2005.  These observations demonstrated that short GRBs are cosmological in origin, reside in both star forming and elliptical galaxies, are not associated with supernovae, and span a wide isotropic-equivalent energy range of $\\sim 10^{48}-10^{52}$ erg. However, a fundamental question remains unanswered: What are the progenitors of short GRBs?  The most popular theoretical model invokes the coalescence of compact object binaries with neutron star and/or black hole constituents.  However, additional possibilities exist, including magnetars formed through prompt channels (massive star core-collapse) and delayed channels (binary white dwarf mergers, white dwarf accretion-induced collapse), or accretion-induced collapse of neutron stars.  In this review I summarize our current knowledge of the galactic and sub-galactic environments of short GRBs, and use these observations to draw inferences about the progenitor population. The most crucial results are: (i) some short GRBs explode in dead elliptical galaxies; (ii) the majority of short GRBs occur in star forming galaxies; (iii) the star forming hosts of short GRBs are distinct from those of long GRBs, and instead appear to be drawn from the general field galaxy population; (iv) the physical offsets of short GRBs relative to their host galaxy centers are significantly larger than for long GRBs; (v) there is tentative evidence for large offsets from short GRBs with optical afterglows and no coincident hosts; (vi) the observed offset distribution is in good agreement with predictions for NS-NS binary mergers; and (vii) short GRBs trace under-luminous locations within their hosts, but appear to be more closely correlated with the rest-frame optical light (old stars) than the UV light (young massive stars).  Taken together, these observations suggest that short GRB progenitors belong to an old stellar population with a wide age distribution, and generally track stellar mass.  These results are fully consistent with NS-NS binary mergers and rule out a dominant population of prompt magnetars. However, a partial contribution from delayed magnetar formation or accretion-induced collapse is also consistent with the data. ",
        "introduction": "\\label{sec:intro}  Gamma-ray bursts (GRBs) are short, intense and non-repeating flashes of $\\gamma$-ray radiation originating at cosmological distances. While GRBs exhibit a broad diversity in their prompt $\\gamma$-ray emission (e.g., duration, spectral shape, peak energy, brightness), they can still be divided into two basic categories: short-duration and long-duration with a separation at about 2 sec \\citep{kmf+93}. The short GRBs have durations as short as $\\sim 10$ msec, while the long events extend to hundreds of seconds.  In addition, short GRBs tend to exhibit harder $\\gamma$-ray spectra than the long-duration events, and generally have a lower fluences.  The basic bimodality of GRB durations provided an early clue that the progenitors of the two classes are likely to be distinct.  Within the broad range of possible scenarios, two popular models have emerged: The collapse of rapidly-rotating massive stars (``collapsars''; \\citep{mw99}) in the case of long GRBs, and the coalescence of compact object binaries (with neutron star and/or black hole constituents -- NS-NS/NS-BH; \\citep{elp+89,pac91,npp92}) in the case of short GRBs. The key attractions of this mapping are the potential for a large energy release from both progenitor classes, and the expected typical timescale for each progenitor: A free-fall timescale of $t_{ff}\\approx 30\\,{\\rm s}\\,(M/10\\,M_\\odot)^{-1/2}(R/10^{10}\\,{\\rm cm})^{3/2}$ for collapsars, and a dynamical timescale of milliseconds for the compact merger remnants of neutron stars and black holes.  However, other progenitor systems have also been proposed for short GRBs, for example magnetars, thought to be the power source behind soft $\\gamma$-ray repeaters \\citep{td95}, accretion-induced collapse (AIC) of neutron stars \\citep{qwc+98}, and delayed magnetar formation through binary white dwarf mergers or white dwarf AIC \\citep{lwc+06,mqt08}.  Until a decade ago, the distances, energy scale, geometry, environments, and progenitors of GRBs, as well as the relation between the two burst classes, remained uncertain due to the lack of precise positions.  The discovery of long-wavelength, long-lived ``afterglows'' from long GRBs in 1997 provided the first glimpse at these properties \\citep{cfh+97,fkn+97,vgg+97}.  Indeed, the sub-arcsecond positions enabled by long GRB afterglow detections demonstrated a cosmological origin \\citep{mdk+97}, an energy scale of $\\sim 10^{51}$ erg \\citep{fks+01,bkf03,bkp+03,bfk03}, significant collimation ($\\sim 10^\\circ$ jets; \\citep{hbf+99,sgk+99}), and direct evidence for relativistic expansion \\citep{tfb+04}.  Intense afterglow and host galaxy observations also linked long GRBs with the deaths of massive stars, mainly through their exclusive location in star forming galaxies (e.g., \\citep{bdk+98,dkb+98,ftm+99}), their strong correlation with the rest-frame ultraviolet (UV) light of their hosts \\citep{bkd02,fls+06}, and their association with Type Ic core-collapse supernovae \\citep{hsm+03,smg+03}.  The afterglows of short GRBs were discovered only in 2005.  Prior to that point only a few afterglow searches were possible due to the relative faintness of the $\\gamma$-ray emission, and hence a low event rate and large and delayed error circles \\citep{hbc+02}.  In retrospect, these searches were woefully inadequate, reaching only about 21 mag at $\\delta t\\approx 1$ day in the optical and $\\sim 0.5$ mJy at $\\delta t\\approx {\\rm few}$ days in the radio.  The launch of NASA's {\\it Swift} satellite in late 2004, provided the first chance for rapid and accurate positions for short GRBs, and indeed led to the detection of the first X-ray \\citep{gso+05,bpp+06}, optical \\citep{ffp+05,hwf+05}, and radio \\citep{bpc+05} afterglows in 2005.  As in the case of long GRBs, the determination of accurate positions revolutionized the study of short GRBs.  First, it led to an association with galaxies at cosmological distances (e.g., \\citep{bpc+05,ffp+05,hwf+05,bpp+06,pbc+06,bfp+07}) and hence an energy scale of $\\sim 10^{49}-10^{52}$ erg (assuming isotropy) \\citep{ber07,nak07}.  Second, it led to the association of some events with elliptical galaxies pointing to an old progenitor population \\citep{bpc+05,gso+05,bpp+06}.  Third, it demonstrated that the afterglow emission is similar to that of long GRBs, albeit with a generally lower luminosity \\citep{bpc+05}.  Finally, it provided a rough estimate of the short GRB event rate \\citep{ngf06}.  {\\it Despite these fundamental results the progenitors of short GRBs remain unidentified at the present.}  The key observational test of the NS-NS/NS-BH merger model, the detection of coincident gravitational waves, is at least several years away.  Similarly, theoretical predictions of early optical/UV emission from an accompanying ``mini-supernova'' \\citep{lp98}, caused by the ejection of radioactive material from the merging system, are highly uncertain, and even the most optimistic predictions lead to a faint and rapidly-fading signal that may be challenging to detect.  Thus, the most promising avenue for progress at the present comes from statistical studies of the environments of short GRBs, both on galactic and sub-galactic scales (e.g., \\citep{pbc+06,ber09,fbf10}). These studies benefit from many of the same techniques that linked long GRBs with the death of massive stars, and allow for comparison with theoretical predictions.  In this review, I present the current state of our knowledge about the redshift distribution of short GRBs, the demographics and detailed properties of their host galaxies, and their locations within their hosts.  The structure of this review is as follows.  In \\S\\ref{sec:disc} I summarize the discovery of short GRB afterglows, and the subsequent identifications of their host galaxies and redshifts; the detailed properties of the hosts (luminosities, metallicities, star formation rates, masses, stellar population ages) are discussed in \\S\\ref{sec:hosts} and \\S\\ref{sec:hosts2}; in \\S\\ref{sec:hst} I discuss the sub-galactic environments of short GRBs, utilizing mainly high-resolution {\\it Hubble Space Telescope} observations; I discuss the possibility of large progenitor offsets (due to kicks or globular cluster origin) in \\S\\ref{sec:kicks}; and finally, in \\S\\ref{sec:conc} I use these results to place constraints on the progenitor population. ",
        "conclusions": "While the sample of short GRBs with afterglow positions is still significantly smaller than that of long GRBs, we have made significant progress in understanding their galactic and sub-galactic environments.  The results of host galaxy imaging and spectroscopy, including high-resolution imaging with {\\it HST}, point to an association of short GRBs with old stellar populations within a range of normal star forming and elliptical galaxies.  In nearly every respect (star formation rates, metallicities, sizes, offsets, light distribution) the environments of short GRBs are distinct from those of long GRBs, indicating that their progenitors are not related to a young progenitor population.  As I showed through the study of short GRB offsets, host galaxy demographics, and the redshift distribution, the current observations are fully consistent with NS-NS/NS-BH binary mergers.  However, a partial contribution from other (mainly old) progenitor channels (e.g., WD-WD mergers leading to magnetar formation, WD/NS AIC) cannot be ruled out at the present.  Currently, we do not have conclusive evidence for significant progenitor kicks, which are only expected in the coalescence model.  Still, a few events with sub-arcsecond optical positions do not directly coincide with bright host galaxies, and yet reside within tens of kpc from bright, low-redshift galaxies.  This may be suggestive of progenitor kicks, but it is also possible that these bursts are associated with fainter hosts (likely at higher redshift) with marginal offsets.  With continued vigilance, and a short GRB discovery rate of about 1 event per month, we are likely to gain further insight into the nature of short GRB progenitors in the next few years, possibly with the first detections (or significant limits) of gravitational waves.  As argued in this review, host galaxy observations of existing and future events will play a central role in our on-going quest to determine the identity of short GRB progenitors."
    },
    "1005/1005.1890_arXiv.txt": {
        "abstract": "s{ The discovery of accelerated expansion using supernova surveys has been one of the most surprising discoveries in cosmology in the past ten years. Present and future surveys, among which SNLS, JDEM or LSST, are based on samples of a few hundreds to a million supernovae. The measurement of their spectroscopic redshifts to investigate dark energy properties is already by far the limiting aspect of such surveys. In this paper, I will discuss and illustrate with SNLS data an approach based solely on photometry to both select supernova candidates and determine their redshift.} ",
        "introduction": "From 2003 to 2008, the Supernova Legacy Survey (SNLS) collected data with the MegaCam imager, a 1 square degree array of 36 CCD with a total of 340 million pixels, over four 1-square degree fields. The data were obtained in a rolling search mode, with a typical time sampling of one point every three to four nights, in four different visible frequency bands $g_M$, $r_M$, $i_M$ and $z_M$. The instrument and scanning strategy were designed specifically for the detection of Type Ia supernovae (SNIa) in the redshift range between $0.2$ and $1.0$.  The standard SNIa selection for the cosmological analyses  in SNLS~\\cite{A06} are based on real-time detection and follow-up spectroscopy (thereafter RTA for Real Time Analysis). Despite the significant amount of time allocated to SNLS for the spectroscopy of its candidates, a spectrum could be obtained for roughly half of them only. This justified additional studies based solely on photometry.  The SNLS photometric analyses described here are independent of the RTA, based on different data processing and selection. As such, they can provide a cross-check to estimate possible biases of the standard scenario. They can also be used for studies requiring larger sets of SNIa, as well as for the selection of other transients. For instance, a photometric analysis of the SNLS data has already led to the identification of a sample of 117 core-collapse supernovae (SNCC), from which was derived the most precise measurement to-date of the rate of SNCC at a mean redshift $z\\sim 0.3$ (Bazin et al., 2009~\\cite{bazin09}).  These proceedings present two major steps towards the use of SNIa from pure photometric studies, i.e. without requiring any spectroscopic information, to derive cosmological constraints. The photometric identification of SNIa, summarizing work which can be found in Bazin et al. (2010) in prep., is described in section~\\ref{sec:identification}. The determination of the photometric redshift of SNIa (see Palanque-Delabrouille et al., 2010~\\cite{palanque} for details) is presented in section~\\ref{sec:redshift}. ",
        "conclusions": ""
    },
    "1005/1005.1774_arXiv.txt": {
        "abstract": "{It is well known that the tilt angles of active regions increase with their latitude (Joy's law). It has never been checked before, however, whether the average tilt angles change from one cycle to another. Flux transport models show the importance of tilt angles for the reversal and build up of magnetic flux at the poles which is, in turn, correlated with the strength of the next cycle.} {Here we analyse time series of tilt angle measurements and look for a possible relationship of the tilt angles with other solar cycle parameters, in order to glean information on the solar dynamo and to estimate their potential for predictions of solar activity.} {We employ tilt angle data from Mount Wilson and Kodaikanal observatories covering solar cycles 15 to 21. We analyse the latitudinal distribution of the tilt angles (Joy's law), their variation from cycle to cycle and their relationship to other solar cycle parameters, such as the strength (or total area covered by sunspots in a cycle), amplitude and length.} {The two main results of our analysis are: 1. We find an anti-correlation between the mean normalized tilt angle of a given cycle and the strength (or amplitude) of that cycle, with a correlation coefficient of $r_{c}=-0.95$ (99.9\\% confidence level) and $r_{c}=-0.93$ (99.76\\% confidence level) for Mount Wilson and Kodaikanal data, respectively. 2. The product of the cycle averaged tilt angle and the strength of the same cycle displays a significant correlation with the strength of the next cycle ($r_{c}=0.65$ at 89\\% confidence level and $r_{c}=0.70$ at 92\\% confidence level for Mount Wilson and Kodaikanal data, respectively). An even better correlation is obtained between the source term of the poloidal flux in Babcock-Leighton-type dynamos (which contains the tilt angle) and the amplitude of the next cycle. Further results are: We confirm the linear relationship (Joy's law) between the tilt angle and latitude with slopes of 0.26 and 0.28 for Mount Wilson and Kodaikanal data, respectively. In addition, we obtain good positive correlations between the normalized area weighted tilt angle and the length of the following cycle, whereas the strength or the amplitude of the next cycle do not appear to be correlated to the tilt angles of the current cycle alone.} {The results of this study indicate that in combination with the cycle strength, the active region tilt angles play an important role in building up the polar fields at cycle minimum.} ",
        "introduction": "Solar cycles differ from each other, showing different lengths, amplitudes and strengths. Understanding the cause of such variations and, ideally, reproducing them is one of the aims of dynamo theory.  Magnetic flux transport dynamo models of the Sun's global magnetic field have been shown to reproduce fairly well the amplitude and duration, among other characteristics, of the solar cycle for at least a few cycles \\citep{charbonneau05, charbonneau07, dikpati06}. Some of the key ingredients of such models include differential rotation, meridional flow, latitude distribution of sunspots, latitudinal drift and a systematic tilt angle of the bipolar groups (Joy's law). These ingredients together explain the polarity reversal of the magnetic field at the poles every $\\sim11$ years. Due to differential rotation the magnetic field lines are wound up around the Sun's rotation axis and when this field is strong enough it becomes buoyant and rises to the surface as sunspots \\citep{babcock61,dikpati08}. The magnetic flux from the sunspots is carried by the meridional flow to the poles, finally causing the reversal. It was already proposed by \\citet{leighton69} that for the reversal to occur there must be cancellation between the leading portions of spots on opposite hemispheres through the slight tilt of the bipolar regions. In this way a greater fraction of the following polarity flux reaches the poles.  \\citet{schrijver02} tested the hypothesis where the polar magnetic field on the Sun is determined by the accumulation of field transported poleward from sunspots at lower latitudes as a consequence of the tilt in the bipoles. Their model was not able to reproduce the polar field measurements of the past years if only the rate at which sunspots emerge is varied from one cycle to another. Furthermore, \\citet{wang02} included a cycle to cycle variable meridional flow in order to achieve agreement with other observations, demostrating that this variable meridional flow could serve as a regulator of the polarity reversal process. Flux-transport simulations also showed that the strength of the polar fields, which feed the dynamo and help determine the strength of the next cycle, is sensitive to the average tilt angle of the active regions of the previous cycle \\citep{baumann04}. Here we investigate whether there is a variation of the cycle averaged tilt angle and Joy's law from cycle to cycle and whether there is a relationship between the tilt angles and the strength, i.e. the activity level, of the following cycle.  Previous studies of the sunspot tilt angles have mainly focused on their relationship with other spot parameters such as magnetic flux, drift motions, rotation, area, polarity separation and cycle phase among others \\citep{howard91-a, howard96-a, sivaraman07}. Variations from one activity cycle to the next, however, have never been explored. This could shed some light on the mechanism by which the magnetic field of active regions is transported to the poles and may thus have the potential for forecasting future solar activity.\\\\  Prediction of future solar activity is one of the main challenges in solar physics and is not only of scientific importance but potentially helps to make predictions about changes to our natural environment that can affect our lives, e.g. space weather and Earth's climate. Most of the present day predictions are based on statistical analyses of solar activity in the past \\citep{hathaway99,hathaway09}. A more physics-based approach is offered by models of the evolution of the Sun's magnetic field, although recent dynamo computations have given controversial results for cycle 24 \\citep{dikpati06, choudhuri07,jiang07}.  The paper is structured as follows. In Sect. 2 we describe the data, the method and some tests. Section 3 presents the results, which are then discussed in Sect. 4. In Sect. 5 we present our main conclusions. ",
        "conclusions": "We have analyzed the sunspot data from Mount Wilson (MW) and Kodaikanal (KK) observatories in order to study Joy's law, the variation of sunspot group tilt angles from cycle to cycle and the relationship of this variation with 3 solar cycle parameters: strength, amplitude and length. The correlations found from the analysis are listed in Tables 1, 2, 3 and 4. From the analysis we highlight the following:\\\\  (1) A linear fit to Joy's law gives $\\alpha=(0.26\\pm0.05)\\lambda$ for the MW and $\\alpha=(0.28\\pm0.06)\\lambda$ for the KK data sets. Here $\\alpha$ is the tilt angle and $\\lambda$ the latitude, both expressed in degrees.  (2) The mean tilt angle changes from cycle to cycle (Fig. 2a and Table 1). The range of values exceeds the uncertainties in the cycle-averaged tilt angles.  (3) A negative correlation between the normalized tilt angle, or $\\langle\\alpha\\rangle/\\langle\\lambda\\rangle$, and the strength of the same cycle is found ($r_{c}=-0.95$ and $r_{c}=-0.93$ for MW and KK data sets, respectively).  (4) We also find a negative correlation between the latitude normalized area weighted tilt angle ($\\langle\\alpha_{w}\\rangle/\\langle\\lambda\\rangle$) and the length of the next cycle ($r_{c}=-0.67$ and $r_{c}=-0.61$ for MW and KK data sets, respectively).  (5) Finally, we discovered a positive correlation between the strength of one cycle multiplied by its mean area- and latitude- weighted tilt angle, $(S\\langle\\alpha_{w}\\rangle/\\langle\\lambda\\rangle)_{i}$, and the strength of the  \\textit{next} cycle, $S_{i+1}$, ($r_{c}=0.65$ and $r_{c}=0.70$ for MW and KK data sets, respectively). Higher correlation coefficients are obtained between a tilt-angle based expression obtained through guidance from Babcock-Leighton type dynamo models and the amplitude of the next cycle ($r_{c}=0.79$ and $r_{c}=0.78$ for MW and KK, respectively).\\\\  These results show the importance of the tilt angle of sunspot groups for both the prediction of solar activity and the understanding of the physics behind the solar dynamo."
    },
    "1005/1005.0258_arXiv.txt": {
        "abstract": "{We provide hot star wind models with radiative force calculated using the solution of comoving frame (CMF) radiative transfer equation. The wind models are calculated for the first stars, O stars, and the central stars of planetary nebulae. We show that without line overlaps and with solely thermal line broadening the pure Sobolev approximation provides a reliable estimate of the radiative force even close to the wind sonic point. Consequently, models with the Sobolev line force provide good approximations to solutions obtained with non-Sobolev transfer. Taking line overlaps into account, the radiative force becomes slightly lower, leading to a decrease in the wind mass-loss rate by roughly 40\\%. Below the sonic point, the CMF line force is significantly lower than the Sobolev one. In the case of pure thermal broadening, this does not influence the mass-loss rate, as the wind mass-loss rate is set in the supersonic part of the wind. However, when additional line broadening is present (e.g., the turbulent one) the region of low CMF line force may extend outwards to the regions where the mass-loss rate is set. This results in a decrease in the wind mass-loss rate. This effect can at least partly explain the low wind mass-loss rates derived from some observational analyses of luminous O stars. ",
        "introduction": "One of the most important galactic populations consists of massive stars, because these stars dominate the spectra of many galaxies and contribute significantly to the mass and momentum input into the interstellar matter. Moreover, massive stars end their active lives in gigantic explosions such as supernovae or even possibly as the progenitors of gamma-ray bursts \\citep[see][] {wooo,yola}, producing huge amounts of heavier elements.   An important property of hot stars that significantly influences their final stages is the stellar wind \\citep[see, e.g.,][for reviews dedicated to hot star winds]{owopo,kkpreh,pulvina}. However, in stellar evolution calculations it is usually unnecessary to know detailed wind properties, but just the amount of mass expelled from the star per unit of time (mass-loss rate) as a function of stellar parameters (e.g., mass, effective temperature, radius, surface metallicity). However, for many hot stars we simply cannot estimate their true mass-loss rate with the precision necessary to calculate evolutionary models. The situation may be less problematic for luminous O stars, for which relatively good agreement between theoretical predictions and observational results seems to exist \\citep[hereafter \\citetalias{nltei}] {pahole,vikolamet,nltei}.  However, the agreement between theoretically predicted mass-loss rates and those derived from observations may be an illusion caused by the neglect of some physical effects in the wind, such as clumping \\citep{bourak,martclump}. As a result, the true mass-loss rates of O stars may be a few times lower than the standard wind theory predicts. This seems to be supported by the observations of \\cite{fuj} of weak wind line profiles of \\ion{P}{v}. Last but not least, the unexpected occurrence of symmetrical X-ray line profiles seems to require relatively low wind mass-loss rates \\citep{kram}.  Possibly unreliable estimates of hot star wind mass-loss rates are also problematic because altough more realistic evolutionary stellar models can be calculated, by including, e.g.,~rotation and magnetic fields, the wind mass-loss rates remain uncertain. Ideally, all observational indicators of mass-loss rate and theoretical models should find and predict similar mass-loss rates. From the observational point of view, more detailed models of line formation in inhomogeneous media may be necessary to obtain reliable line profiles, and consequently also estimate mass-loss rates \\citep{osporcar,chuchcar}.   From the theoretical point of view, disagreement between theory and observations would imply that some of the assumptions used for the hot-star wind modeling are inaccurate. Part of the disagreement may be caused by using incorrect abundances \\citep{spravnez}, although the reason for a disagreement remains mainly unclear. A thorough inspection of all assumptions involved in the modeling is therefore strongly needed. As a first step in this direction, we studied the influence of X-rays on the wind structure of hot stars. It seems that X-rays alone cannot entirely explain the disagreement between theory and observations \\citep{nlteiii} as their influence on wind mass-loss rates is small and they do not strongly affect the ionization fraction of many important ions, especially that of \\ion{P}{v}. On the other hand, the modified ionization equilibrium may affect the X-ray line formation \\citep{lidataky,nlteiii}, and too ling cooling time in the post-shock region \\citep{cobecru,nlteiii} may cause the so-called \"weak wind problem\" \\citep{bourak,martin,linymarko}.  One of the most important approximations in the hot-star wind modeling is the Sobolev approximation \\citep{sobolevprvni,cassob}, which enables us to solve the line radiation transfer analytically. Some studies confirm its applicability in the supersonic part of smooth line-driven winds \\citep{hamki,ppk,pulpren}. However, the applicability of the Sobolev approximation is questionable especially in the regions close to the photosphere bacause of the existence of strong source function gradients \\citep{opsim}. On the other hand, some models avoid using the Sobolev approximation and use only the comoving-frame (hereafter CMF) method of solving the radiative transfer equation \\citep[e.g.,][]{graham}.  We decided to test the applicability of the Sobolev approximation and include the CMF solution of the radiative transfer equation in our wind models. In this first paper of a series, we describe our method, and study the applicability of the Sobolev approximation using models neglecting continuum opacity. ",
        "conclusions": "We have presented hot star wind models in which the radiative force is calculated using the solution of the comoving frame (CMF) radiative transfer equation. The wind models were calculated for three different groups of stellar parameters (corresponding to evolved first stars, O stars, and the central stars of planetary nebulae) to compare the CMF and Sobolev radiative forces for a broader range of stellar parameters.  The comparison of the CMF radiative force with an approximate one calculated by assuming the Sobolev approximation showed that the Sobolev line force is slightly higher due to the neglect of line overlaps. Thus, the mass-loss rate of wind models that include the Sobolev line force is on average a factor of about $1.7$ higher than a more realistic one calculated using CMF wind models. However, we note that the simple Sobolev approximation applied here for reference can be improved to account for line overlaps and continuum absorption \\citep{olson,humrsrybou,sobdif,sobreci}. We emphasize that modern hot star wind models include line overlaps \\citep[e.g.,][]{vikolamet,pahole,graham}.  Without line overlaps in the case of purely thermal line broadening, the Sobolev approximation provides a reliable estimate of the radiative force even close to the wind sonic point. The CAK model therefore provides a good approximation for a solution obtained with non-Sobolev transfer. Below the sonic point, the CMF line force is significantly lower than the Sobolev one partly because of the strong gradients in the source function. This does not influence the mass-loss rate, as the wind mass-loss rate is set in the supersonic part of the wind below the critical point. However, when additional line broadening is present (e.g., the turbulent one) then the region of low CMF line force may extend outwards to the regions where the mass-loss rate is set. This results in a significant decrease in the wind mass-loss rate. We note that this is not a shortcoming of the Sobolev approximation because the Sobolev approximation is applicable to velocities higher than the turbulent velocity in this case.  The influence of turbulent line broadening may cause a dependence of the wind mass-loss rate on the atmospheric turbulent motions. Because theoretical models are not yet able to predict the atmospheric turbulent motions in hot stars in detail, we are unable to provide reliable wind mass-loss rate predictions until the theory of the atmospheric turbulence develops considerably \\citep[however see][]{cant,ae}. Nowadays, hot star evolution seems to be a deterministic one depending only on the initial stellar parameters, i.e., mass, metallicity, and rotational rate. However, as the properties of atmospheric turbulent motions seem to be a non-trivial function of stellar parameters, the evolution of hot stars may become less deterministic, becoming instead dependent on free parameters describing the role of surface turbulence."
    },
    "1005/1005.4432.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {A number of theoretical models vie to explain the $\\gamma$-ray emission from active galactic nuclei (AGN). This was a key discovery of EGRET. With its broader energy coverage, higher resolution, wider field of view and greater sensitivity, the Large Area Telescope (LAT) of the \\textsl{Fermi} Gamma-ray Space Telescope is dramatically increasing our knowledge of AGN $\\gamma$-ray emission. However, discriminating between competing theoretical models requires quasi-simultaneous observations across the electromagnetic spectrum. By resolving the powerful parsec-scale relativistic outflows in extragalactic jets and thereby allowing us to measure critical physical properties, Very Long Baseline Interferometry observations are crucial to understanding the physics of extragalactic $\\gamma$-ray objects.} % aims heading (mandatory) {We introduce the TANAMI program (Tracking Active Galactic Nuclei with Austral Milliarcsecond Interferometry) which is monitoring an initial sample of 43 extragalactic jets located south of $-30$ degrees declination at 8.4\\,GHz and 22\\,GHz since 2007. All aspects of the program are discussed. First epoch results at 8.4\\,GHz are presented along with physical parameters derived therefrom.} % methods heading (mandatory) {These observations were made during 2007/2008 using the telescopes of the Australian Long Baseline Array in conjunction with Hartebeesthoek in South Africa. These data were correlated at the Swinburne University correlator.} % results heading (mandatory) {We present first epoch images for 43 sources, some observed for the first time at milliarcsecond resolution. Parameters of these images as well as physical parameters derived from them are also presented and discussed. These and subsequent images from the TANAMI survey are available at \\url{http://pulsar.sternwarte.uni-erlangen.de/tanami/} } % conclusions heading (optional), leave it empty if necessary {We obtain reliable, high dynamic range images of the southern hemisphere AGN.  All the quasars and BL Lac objects in the sample have a single-sided radio morphology. Galaxies are either double-sided, single-sided or irregular. About $28\\%$ of the TANAMI sample has been detected by LAT during its first three months of operations. Initial analysis suggests that when galaxies are excluded, sources detected by LAT have larger opening angles than those not detected by LAT. Brightness temperatures of LAT detections and non-detections seem to have similar distributions. The redshift distributions of the TANAMI sample and sub-samples are similar to those seen for the bright $\\gamma$-ray AGN seen by LAT and EGRET but none of the sources with a redshift above 1.8 have been detected by LAT.} ",
        "introduction": "}  Blazars are a radio-loud, violently variable, high-luminosity, high-polarization subset of Active Galactic Nuclei (AGN) that show luminosity variations across the electromagnetic spectrum. They typically exhibit apparent superluminal motion along the innermost parsecs of their radio jets. Their observed behaviour suggests they are very compact objects with parsec-scale jets oriented close to our line of sight. Despite decades of observations and modeling of the powerful relativistic jets of AGN, many fundamental questions about them remain unanswered. Jet composition, formation, and collimation are not well understood. {\\bf Neither} are the mechanisms responsible for their propagation, radiation, and interaction with their ambient medium.  The discovery by the EGRET detector onboard the Compton Gamma Ray Observatory (CGRO) that blazars can be bright $\\gamma$-ray emitters \\citep{Hartman1992} was a major breakthrough in the study of AGN. Study of the $\\gamma$-ray blazar population suggested that there is a close connection between $\\gamma$-ray and radio emission, spectral changes, outbursts and the ejection of parsec-scale jet components \\citep{Dondi1995}. Many, but not all, of the radio brightest blazars have been detected at $\\gamma$-ray energies. On the other hand, a small number of extragalactic jets which do not belong to the blazar class have been shown to be bright $\\gamma$-ray sources, as well \\citep[e.g., NGC\\,1275 and Cen A;][]{AbdoLBAS2009}. Currently there are a number of models attempting to explain the observed $\\gamma$-ray emission of blazars and other extragalactic jets \\citep[see][for a review]{Bottcher2007}. Simultaneous broadband spectral energy distribution (SED) measurements across the electromagnetic spectrum are required to discriminate between such models.  Very Long Baseline Interferometry (VLBI) monitoring of blazars is a crucial component of this multiwavelength suite of observations as it provides the only direct measure of relativistic motion in AGN jets allowing us to calculate intrinsic jet parameters such as jet speed, Doppler factor, opening and inclination angles. They also provide the possibility of identifying the location and extent of emission regions. Thus VLBI observations constrain numerical jet simulations and provide tests of the relativistic-beam model \\citep[e.g.,][]{Cohen2007} that are not possible with any other observational technique.  The successful launch of the \\textsl{Fermi} Gamma-ray Space Telescope, formerly known as GLAST \\citep[Gamma-ray Large Area Space Telescope;][]{Gehrels1999}, on June 11th, 2008 was a major milestone in the quest to understand the connection between the low- and high-energy sections of blazar SEDs. The Large Area Telescope \\citep[LAT;][]{Atwood2009} instrument on \\textsl{Fermi} has broader energy coverage (20\\,MeV -- 300\\,GeV), higher resolution, wider field of view (over $20\\%$ of the sky) and a sensitivity $\\sim$30 times greater than EGRET. The combination of its wide field of view with a scanning pattern of observation means LAT observes the entire sky every $\\sim$3 hours making LAT capable of monitoring the sky on timescales from hours to years. Its higher sensitivity at higher energies allows LAT to measure high energy cutoffs which elucidate acceleration mechanisms, radiation and magnetic fields at the source.  %The role of parsec-scale radio monitoring of $\\gamma$-ray bright AGN in the age of \\textsl{Fermi} is widely %appreciated and VLBI programs such as MOJAVE \\footnote{\\url{www.physics.purdue.edu/MOJAVE/}} %and VIPS are in place. However, all of these programs use northern hemisphere telescopes that are %unable to observe targets south of $\\sim$$-$40 degrees declination.  The \\textbf{T}racking \\textbf{A}ctive Galactic \\textbf{N}uclei with \\textbf{A}ustral \\textbf{M}illiarcsecond \\textbf{I}nterferometry (TANAMI) program is a parsec-scale radio monitoring program targeting extragalactic jets south of $-$30 degrees declination. It uses the telescopes of the Australian Long Baseline Array \\citep[LBA, e.g.,][]{Ojha2004b}, and other telescopes in South Africa, Antarctica and Chile to monitor an initial sample of 43 sources at approximately 2-month intervals. The observations are typically made at two frequencies, 8.4\\,GHz and 22\\,GHz, in order to calculate the spectral indices for the core as well as bright jet components. TANAMI began observations in November 2007 before the launch of \\textsl{Fermi} so a proper observation cadence for the targets could be determined. To maximize the usage of TANAMI data, images of sources are made available at \\url{http://pulsar.sternwarte.uni-erlangen.de/tanami/} as soon as they are available. The current target list of TANAMI is maintained at \\url{http://pulsar.sternwarte.uni-erlangen.de/tanami/sample/}.  While the immediate driver for TANAMI was the imminent launch of the \\textsl{Fermi} Gamma-ray Space Telescope, these dual-frequency observations of morphology, motion, and other temporal variations of the most poorly studied third of the AGN sky have a number of other applications. One particularly exciting prospect is the possibility of combining $\\gamma$-ray and radio monitoring of parsec-scale jet kinematics (and production) with a sensitive neutrino telescope.  The ANTARES neutrino telescope (see \\url{http://antares.in2p3.fr}), which is now fully operational and the KM3NeT detector (which has a projected completion date of 2011, see \\url{http://www.km3net.org}) are being designed to detect neutrino point sources\\footnote{See KM3NeT Conceptual Design Report and references therein at \\url{http://www.km3net.org/CDR/CDR-KM3NeT.pdf}}.  Extragalactic jets are among the most promising candidates to be neutrino point sources \\citep{Waxman2007}. Combined TANAMI, \\textsl{Fermi} and ANTARES data will be used to search for neutrino signals correlated with $\\gamma$-ray flares and epochs of jet production. %We will model the expected integral neutrino event rate and estimate the frequency of such events to be seen with KM3NeT. Data filtering and source-stacking techniques to detect neutrinos from blazar jets will also be developed. The results of this work will be the subject of future publications.  In this paper, we describe how we selected our initial source sample (Sect.~\\ref{sec:sample}) and explain our observation and data reduction procedures (Sect.~\\ref{sec:observations}).  We then present the first epoch images of 43 target sources at 8.4\\,GHz in Sect.~\\ref{sec:results} followed by brief notes on individual sources (Sect.~\\ref{sec:notesindiv}). This is followed by a discussion of our results (Sect.~\\ref{sec:discussion}) and we end with our conclusions (Sect.~\\ref{sec:conclusions}). Throughout the paper we use the cosmology $H_0$=73\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{m}$=0.27, $\\Omega_{\\Lambda}$=0.73 where the symbols have their traditional meanings. ",
        "conclusions": "}  Our first epoch 8.4\\,GHz TANAMI images show that the Australian Long Baseline Array and associated telescopes provide high quality images of $\\gamma$-ray blazars that can be used to study the physics of blazars. In many cases these images represent a substantial improvement on published work. The addition of telescopes in Antarctica and Chile is expected to compensate for the loss of the Hartebeesthoek telescope in South Africa and to improve the current $(u,v)$-coverage.  The TANAMI sample has been defined as a hybrid radio and $\\gamma$-ray selected sample of AGN south of $\\delta = -30^{\\circ}$. Of this sample, $84\\%$ of sources show a one-sided morphology, $12\\%$ (all galaxies) are double sided, while one optically unidentified source is compact and one galaxy has a irregular morphology. Of these quasars and BL Lacs have similar morphologies, all being single-sided. The ten galaxies in the sample include five double sided objects, one irregular object and four single sided objects.  About 28\\,\\% of TANAMI sources have been detected by the \\textsl{Fermi} LAT after three months of observations. When galaxies are excluded, initial analysis shows that $78\\%$ of sources detected by the LAT have opening angles $> 30$ degrees compared to just $27\\%$ of non-LBAS sources. This suggests that $\\gamma$-ray bright jets are pointed closer to the line of sight than $\\gamma$-ray faint jets, with the observed opening angles appearing geometrically increased. This result should be regarded as preliminary owing to the modest number of sources available for analysis at this time.   The redshift distribution of the BL Lacs and quasars in the TANAMI sample is similar to the distributions seen for the LBAS and EGRET blazars. No difference is seen between the radio- and $\\gamma$-ray selected subsamples. The redshift distributions of EGRET and LAT detected AGN are also similar. However, most galaxies and none of the five most distant and most luminous sources have been detected by the LAT. Galaxies have the lowest luminosities, BLLacs are more luminous (than galaxies) as a group while quasars dominate the high end of the luminosity distribution. None of the nine most luminous jets have been detected by the LAT so far.   The high end of the brightness temperature distribution is dominated by quasars with the lower end composed mostly of BL Lacs and galaxies. Thirteen sources have a maximum brightness temperature below the equipartition value, and 29 sources have a value below the inverse Compton limit putting about a third of the values above this limit.  The median value is near $3\\times10^{11}$K with the maximum value exceeding $10^{14}$ K. There does not appear to be any significant difference in the brightness temperature distributions of LAT detections and non-detections. Many of the sources with very high brightness temperatures have not yet been detected by the LAT.  The initial results presented here will be augmented by 22\\,GHz images as well as multi-epoch data in future papers to address our scientific questions including those that require spectral and kinematic information. Of particular, interest will be TANAMI epochs observed since the launch of \\textsl{Fermi} satellite thus providing near-simultaneous data at radio, $\\gamma$-ray as well as at intermediate wavelengths from coordinated observations with other instruments. As discussed in Sect.~\\ref{sec:intro}, TANAMI and \\textsl{Fermi} data will also be used with neutrino data from ANTARES and KM3NET for the identification and study of neutrino point sources.    %__________________________________________________________________"
    },
    "1005/1005.0634_arXiv.txt": {
        "abstract": "We study the gravitational dynamics in the early inspiral phase of coalescing compact binaries using Non-Relativistic General Relativity (NRGR) - an effective field theory formalism based on the post-newtonian expansion, but which provides a consistent lagrangian framework and a systematic way in which to study binary dynamics and gravitational wave emission. We calculate in this framework the spin-orbit correction to the newtonian potential at 2.5 PN. ",
        "introduction": "\\begin{eqnarray} V_{\\text{single}} & = & \\frac{G \\, m_2}{2 \\, r^2} \\Big( n^k \\, S_1^{ik} \\{ v_1^i \\vec{v_2}^2+2 \\, v_2^i(\\vec{v_2}^2-2 \\, \\vec{v_1}.\\vec{v_2}) \\} \\nonumber \\\\ & & +S_1^{0k} \\{ 4 \\, r \\, a_2^k-2 (v_1^k -2 \\, v_2^k) \\, \\vec{v_2}.\\vec{n}-3 \\, n^k \\, \\vec{v_2}^2+4 \\, n^k \\, \\vec{v_1}.\\vec{v_2} \\} \\Big) \\end{eqnarray} \\begin{figure} \\begin{center} \\begin{fmffile}{linear} \\newenvironment{Z} {\\begin{fmfgraph*}(72.5,32.5) \\fmfleft{i1,i2} \\fmf{vanilla}{i1,v1,o1} \\fmf{dashes,tension=0}{v2,v1} \\fmf{vanilla}{i2,v2,o2} \\fmfright{o1,o2}} {\\end{fmfgraph*}} \\begin{Z} \\fmfv{decor.shape=circle,decor.filled=empty,decor.size=3thick,label=$v^2$,l.d=.07w}{v2} \\fmfv{decor.shape=circle,decor.filled=full,decor.size=1.5thick,label=$v^3$,l.d=.05w}{v1} \\end{Z} \\begin{Z} \\fmfv{decor.shape=circle,decor.filled=empty,decor.size=3thick,label=$v^3$,l.d=.07w}{v2} \\fmfv{decor.shape=circle,decor.filled=full,decor.size=1.5thick,label=$v^2$,l.d=.05w}{v1} \\end{Z} \\begin{Z} \\fmfv{decor.shape=circle,decor.filled=empty,decor.size=3thick,label=$v^4$,l.d=.07w}{v2} \\fmfv{decor.shape=circle,decor.filled=full,decor.size=1.5thick,label=$v^1$,l.d=.05w}{v1} \\end{Z} \\begin{Z} \\fmfv{decor.shape=circle,decor.filled=empty,decor.size=3thick,label=$v^5$,l.d=.07w}{v2} \\fmfv{decor.shape=circle,decor.filled=full,decor.size=1.5thick,label=$v^0$,l.d=.05w}{v1} \\end{Z} \\end{fmffile} \\end{center} \\caption{Diagrams with single-graviton exchange. The blobs represent spin insertions. The leading order spin vertex introduces a $o(v^2)$ correction.} \\end{figure} ",
        "conclusions": ""
    },
    "1005/1005.2407_arXiv.txt": {
        "abstract": "{ {} {To investigate the poorly constrained sub-mm counts and spectral properties of blazars we searched for these in the \\textit{Herschel}-ATLAS (H-ATLAS) science demostration phase (SDP) survey catalog.} {} {We cross-matched 500$\\mu$m sources brighter than 50 mJy with the FIRST radio catalogue. We found two blazars, both previously known. Our study is among the first blind blazar searches at sub-mm wavelengths, i.e., in the spectral regime where little is still known about the blazar SEDs, but where the synchrotron peak of the most luminous blazars is expected to occur. Our early results are consistent with educated extrapolations of lower frequency counts and question indications of substantial spectral curvature downwards and of spectral upturns at mm wavelengths.  One of the two blazars is identified with a  Fermi/LAT $\\gamma$-ray source and a WMAP source. The physical parameters of the two blazars are briefly discussed. } {These observations demonstrate that the H-ATLAS survey will provide key information about the physics of blazars and their contribution to sub-mm counts.}} ",
        "introduction": "Blazars, comprising BL Lac objects and flat-spectrum radio quasars, are a class of active galactic nuclei characterized by remarkable properties: high luminosities at the two extreme ends of the electromagnetic spectrum (i.e., in the radio and $\\gamma$-ray bands), rapid variability, apparent superluminal jet speeds, flat or inverted radio spectrum, and high polarization in the optical waveband. They are interpreted in terms of relativistic effects due to jets propagating at velocities  close to the speed of light almost along the line-of-sight. Their spectral energy distribution (SED) can be accounted for as a combination of synchrotron emission, peaking at a frequency $\\nu_{p}^s$ that varies from $\\sim 10^{12}\\,$Hz to $\\simgt 10^{19}\\,$Hz (Nieppola et al. 2006), and inverse Compton scattering, assumed to be responsible for the $\\gamma$-ray peak.  Padovani \\& Giommi (1995) classified BL Lac objects into 3 subclasses (low, intermediate, and high frequency synchrotron peaked BL Lacs) depending on the value of $\\nu_{p}^s$. Abdo et al. (2009) extended the classification to all blazars, subdivided into low (LSP; $\\nu_{p}^s<10^{14}\\,$Hz), intermediate (ISP; $10^{14}< \\nu_{p}^s < 10^{15}\\,$Hz), and high (HSP; $\\nu_{p}^s>10^{15}\\,$Hz) synchrotron peaked blazars. For many years, the blazar selection has relied mostly on either radio or X-ray surveys.  Radio surveys preferentially select LSPs, while X-ray surveys favor the HSPs. The resulting global distribution of $\\nu_{p}^s$ is thus bimodal, although objects peaking at intermediate frequencies have also been found. The high sensitivity and nearly uniform sky coverage of the Fermi Gamma-ray Space Telescope made it a powerful tool for providing large blazar samples (Abdo et al. 2010).  In the present paper, we report on the blazars selected at sub-mm wavelengths in the $16\\,\\hbox{deg}^2$ \\textit{Herschel}-ATLAS (H-ATLAS; Eales et al. 2010) science demonstration phase (SDP) field. The \\textit{Herschel} data cover the poorly explored frequency range close to the synchrotron peak of the most luminous LSPs (Fossati et al. 1998; Padovani et al. 2006); $\\nu_p^s$ carries crucial information about key physical parameters, such as the Lorentz factor $\\gamma_p$ of emitting electrons,  the Doppler factor $\\delta$, and the magnetic field strength $B$ ($\\nu_p^s\\propto \\gamma_p^2\\, \\delta\\, B$).   \\begin{figure} \\centering \\includegraphics[scale=0.45]{14637fg1.eps} \\caption{SED of [HB89] 0906+015 (non-simultaneous data), in the observer's frame. Black dots are data taken from the NASA/IPAC Extragalactic Database (NED), red diamonds refer to H-ATLAS data, red circles are the Fermi/LAT data, red asterisks are the new ATCA data, red squares are the new \\textit{Spitzer} data, red stars are UKIDSS data, red hexagrams are the GMRT data, and green triangles are $3\\sigma$ upper limits derived from IRAS maps. The blue solid/dashed lines  represent the best fit to the synchrotron/inverse Compton part of the SED, omitting the optical/UV fluxes, which are probably dominated by the direct emission from the accretion disk (see text). } \\label{fig:sed1}% \\end{figure}  Evolutionary models (De Zotti et al. 2005) predict $\\simeq 0.15\\,\\hbox{blazars\\ deg}^{-2}$ brighter than $50\\,$mJy at $500\\,\\mu$m, the approximate $5\\sigma$ detection limit of the H-ATLAS survey (Rigby et al. 2010, in preparation), so that the full survey may be expected to yield a sample of $\\simeq 80$ blazars selected, for the first time, at sub-mm wavelengths. This prediction assumes a flat radio spectral index $\\alpha$ ($\\alpha\\simeq 0.1$; $S_\\nu \\propto \\nu^{-\\alpha}$) followed by a parabolic decline close to and above a synchrotron peak frequency that increases with decreasing radio luminosity (Fossati et al. 2008).  However, spectral measurements have provided evidence of a substantial spectral curvature downwards from cm to mm wavelengths (Waldram et al. 2007; Sadler et al. 2008). Vieira et al. (2009) find that their highest significance 150 GHz selected radio sources are consistent with flat spectral behavior (or $\\alpha \\simeq 0.1$) between 5 GHz and 150 GHz but a possible steepening between 150 and 220 GHz. Voss et al. (2006) reported an excess (based however on the detection of only three blazars) above extrapolations from lower frequency counts of the surface density at $S\\ge 10\\,$mJy at 1.2 mm, suggesting that at least a fraction of blazars have inverted (i.e. rising with frequency) spectra at mm wavelengths, perhaps due to variability. The H-ATLAS data will also shed light on this issue.  In this paper, we describe (Sect. \\ref{sec:cand}) the identification of blazars among sources detected in the H-ATLAS SDP field and discuss (Sect. \\ref{sec:SED})  the properties of the 2 blazars found. In Sect. \\ref{sec:concl} we summarize our main conclusions. ",
        "conclusions": "\\label{sec:concl}  The \\textit{Herschel} detection of two blazars in the SDP area is interesting in a number of ways. First of all, it confirms that the H-ATLAS survey can play an important role in determining properties of blazars in a crucially important spectral region. With an area of $550\\,\\hbox{deg}^2$, the full H-ATLAS should yield a sample of $\\simeq 80$ blazars, the first statistically significant sample selected at sub-mm wavelengths. The large area is of crucial importance in this context, because the blazar counts are expected to be relatively flat. The H-ATLAS survey will extend by about one order of magnitude downwards in flux the blazar number counts from the {\\sc Planck} survey, which is predicted to have a (95\\% reliability) detection limit of 430 mJy at 545 GHz (Leach et al. 2008).  The number of detected blazars is consistent with the predictions of the De Zotti et al. (2005) model which yields $\\simeq 0.15\\,\\hbox{blazars\\ deg}^{-2}$ brighter than $50\\,$mJy at $500\\,\\mu$m, i.e., 2.4 objects in the 16 $\\hbox{deg}^{-2}$ SDP area. An observational estimate of the blazar surface density at 5 GHz was obtained by Padovani et al. (2007; see their Figs. 4 and 6). Taking into account the steepening at frequencies approaching the synchrotron peak, the effective 5 GHz to $500\\,\\mu$m spectral index yielded by the model is $\\simeq 0.2$. The $500\\,\\mu$m flux limit of 50 mJy then corresponds to 130 mJy at 5 GHz; our results are then in very good agreement with the corresponding blazar surface density found by Padovani et al. (2007) and the relatively flat spectral index assumed by the de Zotti et al. model. If this result is confirmed by data of the full H-ATLAS blazar sample, future models will need to explain in addition evidence of a substantial spectral curvature downwards from cm to mm wavelengths (Waldram et al. 2007; Sadler et al. 2008).  The blazar catalog of Massaro et al. (2009) lists three more blazars in the SDP field: PKS0858-004 with $S_{1.4\\rm GHz}=325\\,$mJy, 1RXS J085920.6+0047 with $S_{1.4\\rm GHz}=39\\,$mJy, and 1RXS J085749.8+0135 with $S_{1.4\\rm GHz}=89\\,$mJy. None of them has a plausible counterpart among the H-ATLAS sources. In the case of PKS0858-004, this implies that the effective 1.4 GHz to $500\\,\\mu$m spectral index has to be steeper than 0.3, i.e., steeper than the effective indices of the two detected blazars (0.13 and 0.25), but still within the range observed for these sources.  Based on the best-fit model estimate of the synchrotron peak frequency, $\\nu_p^s$, both [HB89] 0906+015 and PKS 0907+022 fall in the low frequency synchrotron peaked (LSP) category. This suggests that the H-ATLAS sample will set strong constraints on the abundance of blazars with low ($\\simlt 10^{13}\\,$Hz) values of $\\nu_p^s$, which, in the blazar sequence scenario, are those with more powerful jets, more luminous accretion disks and higher black hole masses (Ghisellini et al. 2010a)."
    },
    "1005/1005.5011_arXiv.txt": {
        "abstract": "Correlated imaging through atmospheric turbulence is studied, and the analytical expressions describing turbulence effects on image resolution are derived. Compared with direct imaging, correlated imaging can reduce the influence of turbulence to a certain extent and reconstruct high-resolution images. The result is backed up by numerical simulations, in which turbulence-induced phase perturbations are simulated by random phase screens inserting propagation paths. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.2517_arXiv.txt": {
        "abstract": "{ The molecular gas in the DR21 massive star formation region is known to be affected by the strong UV field from the central star cluster and by a fast outflow creating a bright shock. The relative contribution of both heating mechanisms is the matter of a long debate. } {By better sampling the excitation ladder of various tracers we provide a quantitative distinction between the different heating mechanisms. } { HIFI observations of mid-$J$ transitions of CO and HCO$^+$ isotopes allow us to bridge the gap in excitation energies between observations from the ground, characterizing the cooler gas, and existing ISO LWS spectra, constraining the properties of the hot gas. Comparing the detailed line profiles allows to identify the physical structure of the different components. } { In spite of the known shock-excitation of H$_2$ and the clearly visible strong outflow, we find that the emission of all lines up to $\\ga2$~THz can be explained by purely radiative heating of the material. However, the new \\textit{Herschel/HIFI} observations reveal two types of excitation conditions. We find hot and dense clumps close to the central cluster, probably dynamically affected by the outflow, and a more widespread distribution of cooler, but nevertheless dense, molecular clumps. } {} ",
        "introduction": "DR21 is a deeply embedded \\HII{} region created by the radiation from at least six OB stars \\citep{Roelfsema}. It sits within a ridge of dense molecular material that obscures the \\HII{}-region at optical wavelengths. The embedded cluster drives a violent bipolar outflow in north-east to south-west direction. It is prominent in the  2~$\\mu$m emission of vibrationally excited H$_2$, tracing hot, shocked gas and in the Spitzer 4.5~$\\mu$m channel \\citep{Garden,Davis}. As the cluster is located close to the eastern edge of the molecular ridge, the eastern, blue-shifted outflow expands in a blister-like fountain, while the western, red-shifted outflow is highly collimated \\citep{Lane}.  Spitzer 8~$\\mu$m images reveal spots of bright PAH emission with a size below 10$''$ \\citep{Marston}. They represent the surfaces of high-density, UV irradiated clumps, forming photon-dominated (or photo-dissociation) regions (PDRs), transition zones from ionized and atomic gas to dense molecular gas where physics and chemistry are dominated by UV radiation from young stars \\citep{HT1999,MolSpa}. These hot and dense regions give rise to the emission of PDR tracers, such as HCO$^+$, high-$J$ CO, atomic and ionized carbon, and atomic oxygen.  An overview on the numerous existing observations in DR21 is given by \\citet{Schneider2006} and \\citet{Jakob}. {\\changed In spite of the wealth of data, the heating of the molecular gas is still debated. The bright H$_2$ emission indicates shock heating in the outflow and the wings of the molecular lines prove its dynamical impact, but excitation models} of the observed emission of CO, \\CI{}, \\CII{}, and \\OI{} by \\citet{Lane} and \\citet{Jakob} have shown that the emission of those tracers cannot be explained by shocks, but is consistent with a pure UV heating, i.e., PDR physics.  To {\\changed quantify the heating of the gas,} we analyze spectra taken with the {\\it HIFI} instrument \\citep{HIFI} on board the {\\it Herschel Space Observatory} \\citep{Herschel} during the performance verification campaign. Observations of high-$J$ HCO$^+$ transitions trace hot material, ionized by UV radiation or X-rays \\citep{Sternberg1995}. Hot water lines are produced in shocked gas \\citep{Snell2005}, and observations of CO isotopes around $J=10$ close the gap in the excitation ladder between ground-based observations and existing ISO data, allowing to obtain a full picture of the temperature distribution.  In Sect. 2 we present the observational data. Sect. 3 compares the measured line profiles to distinguish different components based on their velocity distribution. In Sect. 4 we provide a model for the emission, supporting the PDR character of the source, and discuss the results in Sect. 5. ",
        "conclusions": "The double-peak structure of the predicted CO intensities across the excitation ladder reflects the two different UV fields leading to different excitation conditions. No single-parameter ensemble can fit all data, but distinct low temperature and high temperature components are needed. The parameters of the model are in agreement with independent estimates. A UV flux of $10^5$ Draine fields corresponds to a geometrical distance of 0.06~pc, i.e. 7$''$, from the central cluster, matching the size of the PAH emission \\citep{Marston}. Our clump densities match those determined by \\citet{Jakob} for the extended cool gas, but are slightly higher than their hot-gas density ($4\\,10^5$~cm$^{-3}$). In contrast, \\citet{Jones, vdT} find still somewhat higher densities for the hot gas, up to $10^7$~cm$^{-3}$. The hot ensemble mass is close to the 170~$M_\\odot$ derived from early CO 7-6 observations by \\citet{Jaffe1989}. The total mass of the PDR ensembles falls between the mass limits derived by \\citet{Jakob} from dust observations and from line radiative transfer fits.  While the existing ground-based observations provide a very good constraint on the properties of the extended cool gas, and the ISO lines show the total amount of hot gas, it is only the set of new HIFI data that puts the hot and cold distributions well apart from each other in terms of the temperature structure. While \\citet{Jakob} obtained cooling curves with single peaks, the new data for the 10-9 lines of the CO isotopes and the HCO$^+$ transitions force the fit to a bimodal distribution of excitation conditions. When we exclude the {\\it Herschel} data from the model fit, we obtain a parameter set that shows a UV field that is lower by a factor ten for the hot ensemble, i.e., that would imply molecular clumps farther away from the central cluster. For the cold ensemble, the fitted UV field is also somewhat lower, while all other parameters remain similar to those from the full fit. Only with the {\\it Herschel} data, we therefore obtain a parameter set that is consistent with the source geometry.  As the two-ensemble PDR model is able to fit all of the observed lines, we find no evidence for a shock heating of the dense gas. This is in agreement with the analysis of \\citet{Lane}, explicitely excluding a shock origin of the fine-structure lines, but seems to be in contradiction with the analysis of the line profiles in Sect. \\ref{sect_profiles} that shows excited outflow material. We conclude that the material visible in the blue line wing, characterizing the blister outflow, is contained in dense clumps that are accelerated by the outflow, but that are chemically and energetically fully dominated by the UV field and not by the associated shock."
    },
    "1005/1005.0986_arXiv.txt": {
        "abstract": "{% We report on the ground-based follow-up program of spectroscopic and photometric observations of solar-like asteroseismic targets for the Kepler space mission. These stars constitute a large group of more than thousand objects which are the subject of an intensive study of the Kepler Asteroseismic Science Consortium Working Group 1 (KASC WG-1). The main goal of this coordinated research is the determination of the fundamental stellar atmospheric parameters, which are used for the computing of their asteroseismic models, as well as for the verification of the Kepler Input Catalogue (KIC).} ",
        "introduction": "The Kepler Asteroseismic Science Consortium KASC\\footnote{Kepler Asteroseismic Science Consortium (KASC) is a group of collaborating scientists established to accomplish the activities of the Kepler Asteroseismic Investigation (KAI), represented by Ronald Gilliland (see http://astro.phys.au.dk/KASC).} has been established with the aim of making a full exploitation of the Kepler time-series space photometry. The main goal of KASC is the study of stellar interiors by means of asteroseismic methods.  The largest group of stars selected for asteroseismic targets for Kepler are those which are expected to show solar-like oscillations. Unfortunately, the Kepler data do not provide information on the effective temperature $(T_{\\rm eff})$ surface gra\\-vi\\-ty, $(\\log g)$, metallicity, and projected rotational velocity, $(v\\sin i)$, of these stars, which are crucial data for asteroseismic modeling. Only a small fraction of Kepler asteroseismic targets has been studied in the literature. The effective temperature has been derived either from spectroscopy or photometry only for 15\\% of the solar-like Kepler asteroseismic targets, the surface gravity, for 10\\%, and 5\\% have a measured metallicity.  The Kepler Input Catalog\\footnote{http://archive.stsci.edu/kepler/kepler\\_fov/search.php} (KIC) (Latham et al.\\ 2005) provides values of $T_{\\rm eff}$, $\\log g$ and [Fe/H] derived from Sloan photometry for many Kepler targets. However, the uncertainties of $T_{\\rm eff}$, $\\log g$ and [Fe/H] in the KIC are too large for asteroseismic modeling since they reach $\\pm 200$\\,K in $T_{\\rm eff}$, and $\\pm 0.5$\\,dex, both in $\\log g$ and [Fe/H].  In Fig.\\ \\ref{kic} we show the differences between $T_{\\rm eff}$, $\\log g$ and [Fe/H] of the solar-like oscillators in the KIC and in the literature. As can be seen from the figure, the discrepancies can be as high as 1000\\,K, especially for stars hotter than 5500\\,K. The agreement in $\\log g$ and [Fe/H] is slightly better but still poor. Their values agree to within the error bars, but the latter are so large that they do not allow any meaningful comparison.  Moreover, we note that there is a trend in the differences between [Fe/H] in the KIC and in the literature, which exists both for the spectroscopic and photometric derivations, and which shows that the metal-deficient stars are found to be more metal-poor, and the metal-rich stars, more metal-abundant in the KIC than in the literature. ",
        "conclusions": ""
    },
    "1005/1005.2451_arXiv.txt": {
        "abstract": "From optical spectroscopic measurements we determine that the HD 15407 binary system is $\\sim$80 Myr old. The primary, HD 15407A (spectral type F5~V), exhibits strong mid-infrared excess emission indicative of a recent catastrophic collision between rocky planetary embryos or planets in its inner planetary system. Synthesis of all known stars with large quantities of dust in their terrestrial planet zone indicates that for stars of roughly Solar mass this warm dust phenomenon occurs at ages between 30 and 100 Myr. In contrast, for stars of a few Solar masses, the dominant era of the final assembling of rocky planets occurs earlier, between 10 and 30 Myr age. The incidence of the warm dust phenomenon, when compared against models for the formation of rocky terrestrial-like bodies, implies that rocky planet formation in the terrestrial planet zone around Sun-like stars is common. ",
        "introduction": "Investigating the formation of terrestrial planets outside of the Solar System can give insight into the origin of the Earth-Moon system and whether similar planetary system architectures form around other stars. The study of extra-solar terrestrial planet formation must currently be carried out through indirect searches for the dusty aftermath of formation events. Since the dust is located in the terrestrial planet zone, it will have temperatures of $\\sim$300 K and will emit most of its thermal radiation at $\\sim$10 $\\mu$m. Stars actively undergoing terrestrial planet formation will have substantial amounts of warm dust orbiting in their terrestrial planet zones and will thus show mid-infrared emission in excess of what one would expect from the star alone. Properly establishing the age of such mid-infrared excess systems is necessary for comparison with expectations from planet formation theories and simulations.  We were alerted to excess mid-infrared emission around the star HD 15407A through a paper entitled ``Detection of silica dust in the debris disk around HD 15407'' presented at an October 2009 Spitzer meeting by H. Fujiwara (E. Becklin, personal communication 2009). At this meeting, the age for HD 15407A was suggested to be $\\sim$2 Gyr. We felt this age could be in error, and initiated our own follow-up observations to independently constrain HD 15407A's age. Here we report these observations, our revised value for the age of the HD 15407 binary system, and a discussion of terrestrial planet zone collisions around adolescent Solar mass stars. ",
        "conclusions": "Based on Keck HIRES optical spectroscopy, X-ray flux, and Galactic UVW space motions we estimate the age of the HD 15407 binary system to be $\\sim$80 Myr and identify the system as likely belonging to the AB Dor moving group. Collecting the full sample of known very dusty debris-disk hosting stars, we find that the era of final catastrophic mass accretion for rocky planets around stars of roughly Solar mass occurs later than around intermediate mass stars. Provided that rocky planet formation in the terrestrial planet zone of most Solar mass stars is accompanied by at least one and possibly a few catastrophic collisions of planet-size objects within a 100 Myr interval, then the detection rate and age range of very dusty Solar-type stars $-$ observed to be about one in 300 for ages between 30 and 100 Myr $-$ can be accounted for. Our conclusion is that existing observations of such dusty stars indicates that most Sun-like stars are accompanied by rocky planets in the terrestrial planet zone and that collisions among such forming planets will occur, as is expected in theoretical simulations of planetary origins."
    },
    "1005/1005.1178_arXiv.txt": {
        "abstract": "{In August 2008, the accreting milli-second X-ray pulsar (AMXP), IGR J00291+5934, underwent an outburst lasting $\\sim$ 100 days, the first since its discovery in 2004.} {We present data from the 2008 double-peaked outburst of IGR J00291+5934 from Faulkes Telescope North, the Isaac Newton Telescope, the Keck Telescope, PAIRITEL, the Westerbork Synthesis Radio Telescope and the Swift, XMM-Newton and RXTE X-ray missions. We study the outburst's evolution at various wavelengths, allowing us to probe accretion physics in this AMXP.} {We study the light curve morphology, presenting the first radio--X--ray Spectral Energy Distributions (SEDs) for this source and the most detailed UV--IR SEDs for any outbursting AMXP. We show simple models that attempt to identify the emission mechanisms responsible for the SEDs. We analyse short-timescale optical variability, and compare a medium resolution optical spectrum with those from 2004.} {The outburst morphology is unusual for an AMXP, comprising two peaks, the second containing a `plateau' of $\\sim$ 10 days at maximum brightness within 30 days of the initial activity. This has implications on duty cycles of short-period X-ray transients. The X-ray spectrum can be fitted by a single, hard power-law. We detect optical variability of $\\sim$ 0.05 magnitudes, on timescales of minutes, but find no periodic modulation. In the optical, the SEDs contain a blue component, indicative of an irradiated disc, and a transient near-infrared (NIR) excess. This excess is consistent with a simple model of an optically thick synchrotron jet (as seen in other outbursting AMXPs), however we discuss other potential origins. The optical spectrum shows a double-peaked H$\\alpha$ profile, a diagnostic of an accretion disc, but we do not clearly see other lines (e.g. He I, II) that were reported in 2004.} {Optical/IR observations of AMXPs appear to be excellent for studying the evolution of both the outer accretion disc and the inner jet, and may eventually provide us with tight constraints to model disc-jet coupling in accreting neutron stars.} ",
        "introduction": "\\subsection{Accreting Milli-Second X-ray Pulsars}  Low-Mass X-ray Binaries (LMXBs) comprise a compact primary (neutron star or black hole) and a low-mass (usually K or M class) secondary with the compact object accreting material from its companion by means of Roche lobe overflow. It was postulated in the early 1980s that accreting milli-second X-ray pulsars would form a sub group of neutron star LMXBs (e.g. \\citealp{rad,alp}) and thus prove to be the link between accreting LMXBs and isolated milli-second pulsars. The mechanism believed to be responsible is the spin-up of an old, weakly magnetized, neutron star caused by the accretion of matter and angular momentum from a donor star. This would result in a compact transient X-ray binary (its transient nature demonstrating a period of increased mass transfer rate onto the neutron star) with the neutron star spinning so rapidly that, once the accretion has ceased, it can once again become a rapid radio pulsar (see \\citealp{wij06} and references therein). This was subsequently confirmed observationally with the discovery of pulsations from the transient source SAX J1808.4$-$3658 \\citep{wij98} in 1998 April and the first LMXB to be detected as a milli-second radio pulsar, FIRST J102347.6+003841 \\citep{arc}.  Properties of AMXPs include weaker peak luminosities in outburst than `classical' (i.e. not showing pulsations during outburst) neutron star LMXBs, extremely low-mass companions and faint quiescent X-ray luminosities ($\\le$ 10$^{32}$ erg s$^{-1}$), dominated by a hard power-law component (see \\citealp{wij06,wij} and references therein). For most AMXPs, the duty cycle is a few weeks of outburst every few years. Taking into account the recent (2009 August and September respectively) discoveries of NGC 6440 X-2 \\citep{alt,hei} and IGR J17511$-$305 \\citep{mar09,boz}, there are currently twelve known AMXPs \\citep{wij}. Observationally, these systems have orbital periods of $P_\\mathrm{orb}$ $\\sim$ 40 minutes - 4 hours. Little is known about the optical counterparts of these systems; many have not been detected in the optical/infrared regime, even in outburst (e.g. XTE J1751-305; \\citealp{jon03}). Some LMXBs suffer from a high level of absorption from interstellar dust and/or lie at large distances, and those systems with detections, particularly AMXPs, are often very faint (due to their relatively low-mass donors), especially when in quiescence (see \\citealp{dav09}).  \\subsection{The Accreting Milli-Second X-ray Pulsar IGR J00291+5934}  The accreting millisecond X-ray pulsar IGR J00291+5934 (hereafter ``00291+59'') was first discovered on 2004 December 2 (MJD 53341) in outburst during a routine Galactic Plane Scan of the Cassiopeia region by INTEGRAL \\citep{eck}. It was given the co-ordinates RA 00:29.1, dec +59:34 (J2000.0) with an error radius of 1\\farcm5 (l = 120\\fdg1, b = $-$3\\fdg2). In three successive pointings, the IBIS/ISGRI instrument recorded an average flux of 55 $\\pm$ 5 mCrab (20 -- 60 keV). During its only pointing, JEM-X detected the source with a flux of 23 $\\pm$ 5 mCrab (3 -- 10 keV) \\citep{eck}. On the following day, using the RXTE PCA (Rossi X-ray Timing Explorer Proportional Counter Array), 00291+59 was confirmed as an X-ray pulsar with a spin of 598.88 Hz \\citep{mark04a}, corresponding to a spin period of $\\sim$  1.7 ms. This made it the fastest known AMXP (and the 6th one to be discovered). A sinusoidal frequency modulation of 147.4 minutes (2.46 hours) found in X-ray (by RXTE) was interpreted as the orbital period and no evidence of X-ray eclipses was present \\citep{mark04b}. The fast X-ray variability was stronger and had lower characteristic frequencies than any other neutron star LMXB, being more similar to that of black hole systems \\citep{lin07}.  Within days, optical and infrared counterparts were identified  at RA 00:29:03.06, dec +59:34:19.0 (J2000.0) to an uncertainty of $\\sim$ 0\\farcs5 \\citep{fox,ste}. The source was bright with magnitudes of R $\\sim$ 17.4; J = 16.8 $\\pm$ 0.1; H = 16.8 $\\pm$ 0.3; K = 16.1 $\\pm$ 0.2. Later observations during quiescence gave values of R = 23.1 $\\pm$ 0.1 in 2005 October and K = 19.0 $\\pm$ 0.1 in 2005 January \\citep{tor08a}. During its decay, variability of $\\sim$ 0.3 magnitudes (not associated with the system's orbital period) was noted on timescales of seconds to hours in R-band \\citep{bik}. An optical spectrum of 00291+59 was taken on December 5 by the William Herschel Telescope (WHT) showing `weak evidence' for broad emission features at H$\\alpha$ and He \\small II\\normalsize\\ (4686\\ \\AA) \\citep{roe,tor08a}, both of which are considered diagnostics for outbursting LMXBs. A further spectrum was taken on December 12 (still during outburst) using the 10-metre Keck 1 telescope \\citep{fil,rey} which  detected broad (FWHM = 1200 km s$^{-1}$) emission at H$\\alpha$, H$\\beta$ and He \\small I\\normalsize\\ (6678\\ \\AA) as well as narrow (FWHM = 300 km s$^{-1}$) `very weak' features at He \\small II\\normalsize\\ (4686\\ \\AA). The optical-IR SED during outburst was dominated by thermal emission with a NIR excess thought to be due to synchrotron emission \\citep{tor08a}. Until the 2008 outburst, no further spectra of 00291+59 had been published.  On 2004 December 4, radio observations with the Ryle Telescope, Cambridge yielded a `probable detection' of 1.1 mJy at 15 GHz \\citep{pool}. 12 hours later, at the same frequency, this signal had disappeared below the detection threshold of $\\sim$ 0.6 mJy \\citep{fen04a}, suggesting that this was indeed the transient source. A day later, the Westerbork Synthesis Radio Telescope (WSRT) detected a radio source with mean flux density of 0.250 $\\pm$ 0.035 mJy (at 5 GHz), with evidence of this fading during the 10 hour observation \\citep{fen04a}. Very Large Array (VLA) observations on December 9 gave a detection of 0.17 $\\pm$ 0.05 mJy at 4.86 GHz \\citep{rup}.  A reliable distance to the system is still to be determined. An upper limit of 3.3 kpc based on X-ray absorption has been suggested \\citep{sha}, however \\citet{bur} suggest a value of between 7.4 and 10.7 kpc, based upon the pulsar's spin-down rate. Most recently, an estimate of 2 -- 4 kpc has been made based upon the X-ray luminosity during outburst \\citep{tor08a}. An upper limit on the system's donor has been estimated at 0.16 M$_{\\sun}$, implying it is most likely a hot brown dwarf \\citep{gall}. This was calculated from its mass function and also its similarities with the best studied AMXP, SAX J1808.4$-$3658.  Observations of 00291+59 were taken in quiescence with the 3.6 metre Italian TNG telescope. The data were taken over 3 nights in 2005 August (V and R-bands), September (J, H, K-bands) and November (I-band) \\citep{dav}. Folding of I--band images on the known orbital period yielded a semi-amplitude of 0.28 $\\pm$ 0.17 magnitudes with minimum luminosity at phase 0, consistent with the neutron star being behind the donor. Quiescent observations were also taken using the WHT in Harris-I band (2006 September), which showed short-term flaring of up to 1 magnitude above the average magnitude of 21.83 $\\pm$ 0.18 \\citep{jon}. This data was folded on the known period and, after removing the flares, a semi-amplitude of $\\sim$ 0.06 magnitudes was found, with the system at its brightest at phase 0.34 $\\pm$ 0.03.  Although the 2004 outburst denoted the discovery of 00291+59, it is worth noting that \\citet{rem} retrospectively detected `marginal' (5$\\sigma$) similar outbursts from the source as detected by the RXTE All-Sky Monitor (ASM) in 1998 November (21 $\\pm$  4 mCrab) and 2001 September (18 $\\pm$ 3 mCrab) suggesting an outburst from this source every $\\sim$ 3 years.  \\subsection{The 2008 Outburst} In 2008 August, 00291+59 again became active. This was first detected by the RXTE PCA on August 13 with a flux of $\\sim$ 16 mCrab (2 -- 10 keV). The PCA had been monitoring the source without detection every three days since 2008 May \\citep{chak}. A previous observation on August 10 had not detected the source. The August 13 detection was confirmed by Swift X-Ray Telescope (XRT) observations on August 15 \\citep{mark08}. Optical photometry was obtained using the Wide Field Camera (WFC) on the Isaac Newton Telescope (INT) \\citep{tor08b} and Faulkes Telescope North (FTN) \\citep{rus}, noting an increase of \\textgreater \\ 4 magnitudes from quiescence to values comparable with the 2004 outburst (see Fig. 1). A detection in the UV was made by the Swift UltraViolet/Optical Telescope (UVOT) instrument on August 15 \\citep{mars08} in the UVW2 filter (central wavelength, 1928\\ \\AA). No radio detections were reported, although the WSRT set a 3$\\sigma$ upper limit of 0.16 mJy \\citep{lin} at 4.9 GHz on August 16. After a rapid fade, a second brightening occurred $\\sim$ 1 month later and was detected on September 18, initially in the optical with FTN, subsequently confirmed by Swift X-ray observations \\citep{lewb}. This second outburst was just $\\sim$ 0.2 magnitudes fainter in i$^\\prime$-band than the outburst peak on August 15.  \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics[width=12cm,angle=270]{14382fg1.ps}} \\caption{FT North and INT i$^\\prime$-band light curve of the 2008 outburst. The dotted line represents the approximate i$^\\prime$-band quiescence level.} \\label{1} \\end{figure}  Here, we present well-sampled optical light curves of both 2008 outburst peaks, plus a quasi-simultaneous optical spectrum, radio, NIR, UV and X-ray data. We describe the data collection and reduction in Section 2. We analyse the light curve morphology, the source's optical variability, spectral properties and broadband behaviour in Sections 3 -- 5. We present the evolution of the NIR--optical--UV SEDs (which are the most complete of any AMXP to date) and the first broadband (radio--X--ray) SEDs of 00291+59, and compare this outburst to the 2004 outburst and to other AMXPs in Section 6. Our conclusions are presented in Section 7. ",
        "conclusions": "We present a unique multi-wavelength dataset of the 2008 double-peaked outburst of the AMXP, IGR J00291+5934, from radio to X-ray, with coverage of both outburst peaks (2008 August and September). We show that, optically, the September fade is similar in shape to the previous 2004 outburst, however we note a difference between the 2008 August and September peaks in that the former exhibits a more rapid fade, while the latter displays a distinct plateau phase. This plateau, at near maximum luminosity, lasted $\\sim$ 10 days before the system faded over a longer timescale than in August. The plateau was also evident in the X-ray light curves; for comparison, the X-ray flux faded by $> 4$ orders of magnitude in $< 10$ days during the first outburst peak in August. From Fig. 4, we see that the morphology of the X-ray fade in both peaks, and in 2004, are very similar. The re-brightening of the system within a few weeks in 2008 is unique within AMXPs, and at odds to this source's previous outburst.  We see short-term variability in the optical of $\\sim$ 0.05 magnitudes within the outburst down to a timescale possible of $\\sim$ 240 seconds, but no periodicity, on the previously published orbital period. This suggests changes in the accretion flow (likely caused by the irradiation of the disc) dominate over ellipsoidal modulations due to the companion star when the source is in outburst. In addition, the optical variability during the September plateau was of a smaller amplitude than the August peak, at similar flux levels.  We study the light curve morphology and evolution, presenting the first radio--X-ray Spectral Energy Distributions for this source and the most detailed UV--IR SEDs for any outbursting AMXP. In the optical, the SEDs contain a blue component, which can be fitted by a blackbody, likely from the disc (at maximum luminosity, T $\\sim 3 \\times 10^4$ K) and a transient near-infrared excess. This excess is consistent with a simple model of a synchrotron jet (as seen in other outbursting AMXPs), however we cannot exclude other potential origins such as the presence of a second, cooler (T $\\sim 3.7 \\times 10^3$ K) transient blackbody of uncertain origin. We display UV--IR SEDs for 12 dates during the outburst, and find that the IR excess fades more rapidly than the optical disc flux. Our optical spectrum shows the double-peaked H$\\alpha$ profile of an accretion disc but we do not clearly see other lines (e.g. He \\small I, II\\normalsize) that were observed in 2004.  The lack of a radio detection prevents us drawing many conclusions about the long wavelength part of the SEDs. However, our upper limit is consistent with an optically thick jet (with $\\alpha \\sim$ +0.2) between radio and NIR, as seen in other AMXPs. Our radio--X-ray SEDs suggest that there is a break in the H-band between an optically thick and optically thin jet. Our simple modelling is also able to account for the radio--optical SED from 2004 which also consists of a blue disc and a NIR-excess \\citep{tor08a}. Extending the optically thin jet from infrared--X-ray suggests that the synchrotron jet could account for up to $\\sim$ 10$\\%$ of the observed X-ray emission, although we note that one synchrotron jet model is able to reproduce the observed hard X-ray power law if the break in the jet spectrum is at a higher frequency than the OIR SED data suggest.  The unusual nature of this double-peaked outburst has implications for the study of AMXPs in particular, and neutron star X-ray transients in general. Not only is this source's double-peaked outburst intriguing, it is difficult to understand what physical processes were different to cause the previous outburst in 2004 to follow a more `traditional' single-peaked morphology. Observing such systems can place tighter constraints on the range of duty cycles of neutron star XRB outbursts. For short period systems such as AMXPs, it allows us to study the evolution of both the outer accretion disc and the inner synchrotron jet, providing data for more accurate modelling of disc-jet coupling in neutron star XRBs. This study illustrates the importance of rapid, regular multi-wavelength monitoring of AMXP outbursts, and of regular optical monitoring of X-ray binaries in quiescence. As was the case with the second outburst peak of IGR J00291+5934, new outbursts can often be detected by such monitoring with 2-metre class telescopes such as the Faulkes Telescopes and YALO/SMARTS (e.g. \\citealp{jai,mai07}) before X-ray detection."
    },
    "1005/1005.1664_arXiv.txt": {
        "abstract": "We show that in cases of marginal detections ($\\sim 3\\sigma$), such as that of Baryonic Acoustic Oscillations (BAO) in cosmology,  the often-used Gaussian approximation to the full likelihood is very poor, especially beyond $(\\sim 3\\sigma$). This can radically alter confidence intervals on parameters and implies that one cannot naively extrapolate $1\\sigma$ error bars to $3\\sigma$ and beyond. We propose a simple fitting formula which corrects for this effect in posterior probabilities arising from marginal detections. Alternatively the full likelihood should be used for parameter estimation rather than the Gaussian approximation of a just mean and an error. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.0935_arXiv.txt": {
        "abstract": "Inflation is studied in the context of induced gravity (IG) $\\gamma \\sigma^2 R$, where $R$ is the Ricci scalar, $\\sigma$ a scalar field and $\\gamma$ a dimensionless constant, and diverse symmetry-breaking potentials $V(\\sigma)$ are considered. In particular we compared the predictions for Landau-Ginzburg (LG) and Coleman-Weinberg (CW) type potentials and their possible generalizations with the most recent data. We find that large field inflation generally leads to fewer constraints on the parameters and the shape of the potential whereas small field inflation is more problematic and, if viable, implies more constraints, in particular on the parameter $\\gamma$. We also examined the reheating phase and obtained an accurate analytical solution for the dynamics of inflaton and the Hubble parameter by using a multiple scale analysis (MSA). The solutions were then used to study the average expansion of the Universe, the average equation of state for the scalar field and both the perturbative and resonant decays of the inflaton field. ",
        "introduction": "Many years ago a model for a varying gravitational coupling was introduced \\cite{brans}. The model consisted of a massless scalar field whose inverse was associated with the gravitational coupling. Such a field evolved dynamically in the presence of matter and led to cosmological predictions differing from Einstein Gravity (EG) in that one generally obtained a power-law time dependence for the gravitational coupling. Subsequently it was suggested that the gravitational constant is generated as a one loop effect in some fundamental interaction \\cite{sakharov} or through spontaneous symmetry breaking (for reviews see \\cite{Zee:1980sj,adler}). In order to reduce the strong time dependence in a cosmological setting, which remained after the introduction of matter, a simple globally scale invariant model for induced gravity (IG) involving a scalar field $\\sigma$ and a quartic potential $\\lambda \\sigma^{4}$ was introduced \\cite{CV}. The spontaneous breaking of scale invariance in such a context, either through the presence of a condensate \\cite{zee,Accetta,Sorbo} or quantum effects \\cite{CW,spokoiny}, then led to EG plus a cosmological constant and, on treating matter as a perturbation, a time dependence for the scalar field (gravitational constant) and consistent results \\cite{CV}. A more detailed analysis \\cite{FTV} of such a simple model including both radiation and matter showed that it led to EG plus a cosmological constant as a stable attractor among homogeneous cosmologies and was therefore a viable Dark Energy model for a range of scalar field initial conditions and a positive $\\gamma$ coupling to the Ricci scalar $\\gamma\\sigma^{2}R$. In that earlier study we considered values for the scalar field which were sufficiently close to the spontaneously broken symmetry equilibrium values for the scalar field and compared our results with present values of the cosmological constant and the solar system data. In a later note \\cite{CFTV} we studied the above approach for values of our parameters sufficiently far from the equilibrium values (back in time) for sufficient inflation and the subsequent reheating to take place and examined the compatibility of the predictions for different symmetry breaking potentials with the current data.\\\\ The aim of this paper is to investigate in detail the dynamics of the scalar field and the generation of scalar and tensor perturbations in the early Universe.\\\\ The paper is organized as follows: in section II the formalism and the general equations for the homogeneous dynamics are described. In section III we formalize the problem of perturbations in the IG context and section IV is dedicated to comparing the predictions of diverse symmetry-breaking potentials with the most recent data. In section V we apply a Multiple Scale Analysis (MSA) to the phase of coherent oscillations of the inflaton field at the end of inflation. In section VI we study the perturbative decay of the inflaton field in the IG context and in section VII the resonant decay of the inflaton is discussed. Finally, in section VIII, our conclusions are summarized. ",
        "conclusions": "We have investigated in detail inflation in the IG framework where inflation is driven by the same scalar field as that associated with the observed value of Newton's constant. Inflation in IG can be successful in the SR regime, and leads to a nearly flat spectrum of scalar perturbations and a small tensor to scalar ratio compatibly with observations. The SR conditions are not simply associated with the shape of the inflaton potential since the dynamics of the scalar field strongly depends on its coupling $\\gamma$ to gravity. We have made accurate comparisons with observations for different ``symmetry breaking'' potentials and showed how their shape is relevant for successful predictions of the model. In particular we have seen that the LG and CW type potentials are compatible with the data for large field inflation and for suitable values of the parameters for small field inflation. For potentials involving higher powers of the scalar field (with respect to the previously mentioned cases) agreement with the data is more problematic and if it occurs one obtains constraints on $\\gamma$ for both the large and small field cases.  We have also studied, always in IG, the coherent oscillatory regime and obtained an accurate analytical solution for the dynamics of inflaton and of the Hubble parameter by using a MSA. We employed the solution to first investigate the behavior of the average evolution of the Universe. We found that, independently of the average equation of state of the scalar field, the Universe expands as $\\av{H(t)}\\propto2/(3t)$.\\\\ Further perturbative and resonant reheating have also been studied. We found, in the former case, that the decay of the inflaton into ordinary matter is less efficient for $\\gamma$ large and that the reheating temperature is comparable with the EG value for $\\gamma$ small. We also examined the parametric amplification of the scalar and tensor perturbations and of a generic scalar field non-minimally coupled to gravity and coupled to the inflaton. We note that parametric resonance occurs for scalar perturbations with $H\\ll k/a\\ll\\omega_{0}$ and such perturbations maintain a constant amplitude instead of decaying as the Universe expands. Finally parametric resonance does not occur for gravitational waves whereas amplification may occur for a generic scalar field depending on the parameters of its Lagrangian.  \\appendix"
    },
    "1005/1005.2101_arXiv.txt": {
        "abstract": "We investigate whether the recovery chances of highly spinning waveforms by matched filtering with randomly chosen spinning waveforms generated with the LAL package are improved by a cross-correlation of the simulated output of the L1 and H1 LIGO detectors. We find that a properly defined correlated overlap improves the mass estimates and enhaces the recovery of spin angles. ",
        "introduction": "Recovering a gravitational wave pattern in the noisy detector output is a difficult problem. In this work we present our attempt to recover a gravitational spinning waveform $h_{inj}$ immersed in S5 LIGO-like noise by a modified matched filtering method. We calculate the overlap \\cite{Overlap} between the noisy injection $h_{n,i}$ and spinning templates $h_{template}$ as% \\begin{equation} O\\left[ h_{n,i},h_{template}\\right] =\\frac{\\left\\langle h_{n,i}|h_{template}\\right\\rangle }{\\sqrt{\\left\\langle h_{n,i}|h_{n,i}\\right\\rangle \\left\\langle h_{template}|h_{template}\\right\\rangle }}~,  \\label{O} \\end{equation}% with% \\begin{equation} \\left\\langle h_{1}|h_{2}\\right\\rangle =4\\mathit{Re}\\int_{f_{\\min }}^{f_{\\max }}\\frac{\\widetilde{h}_{1}\\left( f\\right) \\widetilde{h}_{2}^{\\ast }\\left( f\\right) }{S_{n}\\left( f\\right) }df~, \\end{equation}% where tilde and star denote Fourier transform and complex conjugate, respectively, and $S_{n}\\left( f\\right) $ is the power spectral density of the noise. We choose $f_{\\min }=50$Hz and $f_{\\max }=600$Hz, lying in the best sensitivity band of the LIGO detectors, also the post-Newtonian prediction for the waveform is quite accurate there. \\begin{figure}[th] \\begin{center} \\includegraphics[height=7.0cm, angle=270]{injectedhplus.ps} \\hskip 0.5cm % \\includegraphics[height=7.0cm, angle=270]{injectedhcross.ps} \\end{center} \\caption{ The polarization $h_{+}$ (left) and $h_{\\times }$ (right) of the spinning waveforms.} \\label{fig1} \\end{figure}  On theoretical grounds, infinite long data series would be required for exact determination of the power spectrum. In order to achieve stability of the power spectral density of order of 1\\%, we would like to have at least 100 periods of the lowest frequency, therefore a minimal length of the templates of 2 sec was imposed. We compute the above overlap both for the Hanford and Livingston detectors ($O^{H}$ and $O^{L}$, respectively).  We also defined a \\textit{correlated match}, the overlap $O_{corr}^{HL}$, using the correlations of the Hanford and Livingston signals for $h_{n,i}$, $% h_{template}$ and noises. (The correlation of $h_{1}$ and $h_{2}$ in Fourier space is defined as $\\widetilde{H}\\left( f\\right) =\\widetilde{h}_{1}\\left( f\\right) \\widetilde{h}_{2}^{\\ast }\\left( f\\right) $.) \\begin{figure}[th] \\begin{center} \\includegraphics[height=7.0cm, angle=270]{injectedHanford.ps} \\hskip 0.5cm % \\includegraphics[height=7.0cm, angle=270]{injectedLivingstone.ps} \\vskip % 0.2cm \\includegraphics[height=7.0cm, angle=270]{injectednoisyHanford.ps} % \\hskip 0.5cm % \\includegraphics[height=7.0cm,angle=270]{injectednoisyLivingstone.ps} \\end{center} \\caption{ Spinning waveforms at Hanford (top left) and Livingston (top right), mixed in LIGO S5-like noise (bottom).} \\label{fig2} \\end{figure} \\begin{figure}[th] \\begin{center} \\includegraphics[height=7.0cm, angle=270]{unfilterednoise.ps} \\hskip 0.5cm % \\includegraphics[height=7.0cm, angle=270]{filterednoise.ps} \\end{center} \\caption{ Characteristic unfiltered (left) and filtered (right) LIGO noises.} \\label{fig3} \\end{figure}  For the two detectors we also employed the quantity $O^{HL}$, according to Ref. \\cite{SpinSpiral}:% \\begin{equation} O^{HL}=\\frac{\\left\\langle h_{n,i}|h_{template}\\right\\rangle _{H}+\\left\\langle h_{n,i}|h_{template}\\right\\rangle _{L}}{\\sqrt{\\left( \\left\\langle h_{n,i}|h_{n,i}\\right\\rangle _{H}+\\left\\langle h_{n,i}|h_{n,i}\\right\\rangle _{L}\\right) \\left( \\left\\langle h_{template}|h_{template}\\right\\rangle _{H}+\\left\\langle h_{template}|h_{template}\\right\\rangle _{L}\\right) }}~. \\end{equation}% Therefore we could compare four type of matches.  The spinning waveforms, based on Ref. \\cite{BCV2}, were computed by using SpinTaylor Ref. \\cite{SpinTaylor}, while the antenna functions based on Ref. \\cite{Antenna} with the XLALComputeDetAMResponse() function in DetResponse.c under the LAL package Ref. \\cite{SpinTaylor}.  \\begin{table}[t] \\caption{The parameters of the gravitational wave polarizations ($% h_{+},h_{\\times }$) of the injected signal. Masses $m_{i}$, magnitude of the dimensionless spins $\\protect\\chi _{i}$, initial direction of the spin vectors given by $\\cos \\protect\\kappa _{i}$ and $\\protect\\psi _{i}$ in the frame with the line of sight on the $z$-axis, and the initial orbital angular momentum $\\mathbf{L}_{\\mathbf{N}}$ in the $x$-$z$ plane, spanning the angle $\\Theta $, and the distance $d_{L}$ of the source. The initial phase is $0$; the initial time being also fixed to $0$.}$% \\begin{tabular}{c|c|c|c|c|c|c|c|c|c|c} name & $m_{1}(M_{\\odot })$ & $m_{2}(M_{\\odot })$ & $\\chi _{1}$ & $\\chi _{2}$ & $\\cos \\kappa _{1}$ & $\\psi _{1}$ & $\\cos \\kappa _{2}$ & $\\psi _{2}$ & $% \\Theta $ & $d_{L}\\left( Mpc\\right) $ \\\\ \\hline injection & $3.553$ & $3.358$ & $0.983$ & $0.902$ & $0.984$ & $1.109$ & $% 0.978$ & $0.957$ & $1.430$ & $1$% \\end{tabular}% $% \\label{table1} \\end{table} \\begin{table}[t] \\caption{The angles $\\protect\\theta $, $\\protect\\varphi $ and $\\protect\\psi $ (polarization angle) give the relation between a fictitious Earth-centered detector and the source frame. For the actual detector positions further rotation have to be taken into account \\protect\\cite{Antenna}. These angles are also necessary for computing the antenna functions ($F_{+},F_{\\times }$% ). The gravitational signal is $h=h_{+}F_{+}+h_{\\times }F_{\\times }$.} \\begin{center} $% \\begin{tabular}{c|c|c|c} name & $\\varphi $ & $\\theta $ & \\multicolumn{1}{|c}{$\\psi $} \\\\ \\hline injection & $3.657$ & $0.278$ & \\multicolumn{1}{|c}{$0.000$}% \\end{tabular}% $% \\end{center} \\label{table2} \\end{table}  Table \\ref{table1} contains the parameters of the injected waveform. On Fig % \\ref{fig1} the two polarizations of this waveform $h_{+}$ and $h_{\\times }$ are shown. The parameters characterizing the source and the detector orientation, which are necessary for computing the antenna functions $F_{+}$ and $F_{\\times }$ are given in Table \\ref{table2}. The waveforms appearing at the Hanford and Livingston detectors are plotted on Fig \\ref{fig2} (top line), while the bottom line on Fig \\ref{fig2} shows the signals immersed in S5 LIGO-like noise. Some characteristic unfiltered and filtered noises can be seen on Fig \\ref{fig3}.  First we calculated the four type of overlaps between the noisy injection and the injected signal itself, finding $O^{H}\\left[ h_{n,i},h_{injection}% \\right] =0.8392$, $O^{L}\\left[ h_{n,i},h_{injection}\\right] =0.9064$, $O^{HL}% \\left[ h_{n,i},h_{injection}\\right] =0.8662$ and $O_{corr}^{HL}\\left[ h_{n,i},h_{injection}\\right] =0.8646$. Next we defined the following auxiliary quantity for all overlaps% \\begin{equation} \\sigma =\\left\\vert 1-\\frac{O\\left[ h_{n,i},h_{template}\\right] }{O\\left[ h_{n,i},h_{injection}\\right] }\\right\\vert ~, \\end{equation}% vanishing for the injected template. We searched for templates with the various $\\sigma $-s less then $0.1$. \\begin{figure}[th] \\includegraphics[height=16cm, angle=270]{params.ps} \\caption{The parameters of \\ the injected signal (red) and the templates having $\\protect\\sigma <1$ according to $O^{H}$ (green), $O^{L}$ (yellow), $% O^{HL}$ (dark blue) and all of these (light blue) and when the correlated overlap was also taken into account (black). The big dots among the black represent those templates that has the smallest values accordint to the correlated overlap.} \\label{fig4} \\end{figure}  The templates were chosen with parameters in the ranges: masses $m_{i}\\in $ $% 3\\div 10$ M$_{\\odot }$; dimensionless spins $\\chi _{i}\\in 0.7\\div 1$; spin angles $\\cos \\kappa _{i}$ and $\\psi _{i}$ random; distance $d_{L}$, angles $% \\Theta $, $\\varphi $, $\\theta $ and $\\psi $ are fixed identically to the values given in Tables \\ref{table1} and \\ref{table2} for the injection, respectively. We also assumed as known the time of signal arrival to the detector. Therefore we varied two mass and six spin parameters altogether. \\begin{figure}[th] \\includegraphics[height=7.0cm, angle=270]{162197_templateH.ps} \\hskip0.5cm % \\includegraphics[height=7.0cm, angle=270]{162197_templateL.ps} \\vskip0.2cm % \\includegraphics[height=7.0cm, angle=270]{327273_templateH.ps} \\hskip0.5cm % \\includegraphics[height=7.0cm,angle=270]{327273_templateL.ps} \\vskip0.2cm % \\includegraphics[height=7.0cm, angle=270]{980790_templateH.ps} \\hskip0.5cm % \\includegraphics[height=7.0cm, angle=270]{980790_templateL.ps} \\caption{Best three templates (black) as they would appear at Hanford (left) and Livingston (right), compared to the respective injected signals (red).} \\label{fig5} \\end{figure} ",
        "conclusions": ""
    },
    "1005/1005.2930_arXiv.txt": {
        "abstract": "Of the over 450 exoplanets known to date, more than 420 of them have been discovered using radial velocity studies, a method that tells nothing about the inclination of the planet's orbit. Because it is more likely that the companion is a planetary-mass object in a moderate- to high-inclination orbit than a low-mass stellar object in a nearly face-on orbit, the secondary bodies are presumed to be planets. Interferometric observations allow us to inspect the angular diameter fit residuals to calibrated visibilities in order to rule out the possibility of a low-mass stellar companion in a very low-inclination orbit. We used the Center for High Angular Resolution Astronomy (CHARA) Array interferometer to observe 20 exoplanet host stars and considered five potential secondary spectral types: G5 V, K0 V, K5 V, M0 V, and M5 V. If a secondary star is present and is sufficiently bright, the effects of the added light will appear in interferometric observations where the planet will not. All secondary types could be eliminated from consideration for 7 host stars and no secondary stars of any spectral type could be ruled out for 7 more. The remaining 6 host stars showed a range of possible secondary types. ",
        "introduction": "Radial velocity observations of exoplanet systems alone are insufficient to distinguish between intermediate- to high-inclination planetary systems and low-inclination binary star systems. \\citet{2001A&A...371..250S} estimated probability densities of orbital periods and eccentricities for two samples: exoplanet candidates and spectroscopic binary star systems with solar-type primary stars. They found the distributions of the two populations were statistically indistinguishable in the context of orbital elements.  In an earlier study, \\citet{1998A&A...334L..37I} modeled nine exoplanet systems as binary star systems to test if the radial velocity observations could be reproduced by low-mass stellar companions. Although the probability of binary star systems appearing as planetary systems was low -- 0.01 to 4$\\%$ -- the model results described the observations satisfactorily and showed it is possible for a binary star system to mimic an exoplanet system.  While it is unlikely that there are unseen stellar companions with very low inclinations masquerading as exoplanets, the only way to exclude the possibility of a low-inclination star observationally is to study the system at angular scales comparable to the calculated star-planet separation. These separations are on the order of 0.5 to 5.0 milliarcseconds, and the only relevant technique applicable is interferometry, which we present here for 20 exoplanet host stars.  The probability of a system's inclination being in the range $i$ to $i + \\Delta i$ is proportional to the ratio of the surface element of a hemisphere defined by that range and integrated over the azimuth angle (\\emph{$\\Phi$}) to the surface area of the entire hemisphere. The area element for a given range of \\emph{i} is: \\begin{equation} \\mathrm{d} A = \\mathrm{d} i \\times \\mathrm{sin} i \\; \\mathrm{d} \\Phi, \\end{equation} and the probability of a system having a specific range of \\emph{i} is: \\begin{equation} P_{i,i+ \\Delta i} = \\frac{\\int^{2 \\pi}_0 \\int^{i+ \\Delta i}_i \\sin i \\; \\mathrm{d}i \\; \\mathrm{d} \\Phi} {\\int^{2 \\pi}_0 \\int^{\\pi / 2}_0 \\sin i \\; \\mathrm{d} i \\; \\mathrm{d} \\Phi} = \\frac{-2 \\pi \\; \\cos i \\mid^{i+ \\Delta i}_i}{-2 \\pi \\; \\cos i \\mid^{\\pi / 2}_0} = \\cos i - \\cos (i+ \\Delta i). \\end{equation} Therefore the probability of an orbit with an inclination below 45$^\\circ$ is $\\sim$30$\\%$ while the probability of the orbit having an inclination higher than 45$^\\circ$ is $\\sim$70$\\%$. The inclinations necessary for the companion to be stellar in nature are very low, on the order of less than one degree. This means the probability of any given system being a binary star is correspondingly low ($\\sim 10^{-4} \\%$) and it would take much larger sample size than is available to have a good chance of finding a stellar companion masquerading as a planet. When taken as a whole, the chances of finding a stellar companion in this sample are only on the order of $10^{-5}$.  We are not trying to prove that the apparent companion for any given system is a low-mass star instead of a planet, but instead are using our observations to rule out certain types of stellar companions for each star observed. Radial velocity studies alone cannot perform this task, and interferometry is well suited to further our knowledge of exoplanet host stars.  When we are able to be more confident that they are indeed planetary systems and not face-on binary stars, we further characterize the host stars themselves as well as contribute to the statistics that tell us what percentage of stars host planets versus how many have stellar companions.  An important criteria for the detection or non-detection of otherwise unseen secondary stars is the magnitude difference between the known primary and the putative secondary stars. Simulations using a program written by Theo ten Brummelaar that realistically models instrumental and atmospheric noises, as well as observations of pairs of known brightness contrasts, indicate that the Array is sensitive to a magnitude difference in the $K$-band ($\\Delta K$) of 3.0. Therefore, if a second star is present and is not more than $\\sim$3.0 magnitudes fainter than the host star, the effects of the second star will be seen in the interferometric data. It should be noted that a limiting magnitude difference in the $K$-band ($\\Delta K$) of 3.0 is a lower limit, as the true $\\Delta K$ also depends on the absolute brightness of the two stars and could be slightly higher for some systems.  This technique of using interferometric observations to eliminate the possibility of certain types of secondary stars was employed to examine the exoplanet host star 51 Peg (HD~217014) by \\citet{1998ApJ...504L..39B}, whose analysis of Palomar Testbed Interferometer data supported a single-star model for that star. They fit single-star and binary-star models to the data and found that any possible unseen stellar companion would have to have a $K$ magnitude fainter than 7.30 and a mass of less than 0.22$M_\\odot$.  Here we describe our interferometric observations, our method for choosing calibrator stars, and define the role interferometric resolution plays in Section 2. In Section 3, we discuss how the angular diameter fit residuals to calibrated visibilities can help us eliminate certain types of secondary stars, and Section 4 explores the implications of the observations. This paper is follow-on work to an earlier study \\citep{2008ApJ...682..577B}. ",
        "conclusions": "Though the errors for the individual calibrated visibilities listed in Table \\ref{calib_visy} are on the order of 2-20$\\%$, the errors for the angular diameter fits to these visibilities are on the order of 1-6$\\%$ for most stars. This is because the calibrated visibility errors are overestimated. This can be best illustrated by performing the diameter fit to the data. For each $\\theta_{\\rm LD}$ fit, the errors were derived via the reduced $\\chi^2$ minimization method \\citep{2003psa..book.....W,1992nrca.book.....P}: the diameter fit with the lowest $\\chi^2$ was found and the corresponding diameter was the final $\\theta_{\\rm LD}$ for the star. The errors were calculated by finding the diameter at $\\chi^2 + 1$ on either side of the minimum $\\chi^2$ and determining the difference between the $\\chi^2$ diameter and $\\chi^2 +1$ diameter. In calculating the diameter errors in Table~\\ref{host_obs_props}, we adjusted the estimated visibility errors to force the reduced $\\chi^2$ to unity because when this is omitted, the reduced $\\chi^2$ is well under 1.0, indicating we are overestimating the errors in our calibrated visibilities.  In addition to measuring the angular diameters of these stars interferometrically, we also estimated their diameters using two methods to check for discrepancies. We performed SED fits using the method described in Section 2 as well as using the relationship described in \\citet{2004AandA...426..297K} between the ($V-K$) color and log $\\theta_{\\rm LD}$. Table \\ref{ld_sed_vk} lists the results of these calculations and Figure \\ref{diam_compare} plots $\\theta_{\\rm SED}$ and $\\theta_{(V-K)}$ versus $\\theta_{\\rm measured}$. For stars larger than $\\sim$0.7 mas, the errors in the estimated diameters are larger than those for the measured diameters.  In order to characterize the scatter in the diameters of the entire sample, the standard deviation $\\sigma$ of the quantity $| \\theta_{\\rm LD} - \\theta_{\\rm SED} |$ was determined to be $8\\%$, which indicates a fairly good correspondence between the estimated and measured diameters. For comparison purposes, the standard deviation of $| \\theta_{(V-K)} - \\theta_{\\rm SED} |$ was $12\\%$. Four of the 20 stars in the sample have measured diameters that are not within 1-sigma of the diameters estimated using either SED fits or ($V-K$) color. Three of the four (HD~145675, HD~154345, and HD~185269) are among the smallest stars measured and have some of the highest diameter errors in the sample, ranging from 6$\\%$ to 11$\\%$.  The largest outlier is HD~217107, which we measured at 0.70$\\pm$0.01 mas while the diameters from SED fits and ($V-K$) color were 0.52$\\pm$0.02 mas and 0.54$\\pm$0.02 mas, respectively. There are no signs of variability indicated in the literature for the star that would impact the diameters estimated using photometry. While the calibrator was small (0.31$\\pm$0.01 mas), showed no signs of having a stellar companion using speckle \\citep{1987AJ.....93..183M} or in the SED fit, and was used as a photometric comparison star for HD~217107 \\citep{2005ApJ...632..638V}, it could be the cause of the discrepancy in the angular diameter estimates and interferometric measurements. Future observations of HD~217107 using the CHARA Array and different calibrators should help clarify the situation.  The star that showed the most potential for being a binary system instead of a planetary system was the newly-presented HD~164922. The visibility points show a slight sinusoidal pattern, though there are not enough observations to reliably fit the data to a binary star model. Future planned observations using the CHARA Array over a longer time should provide more details on the star.  No secondary spectral types could be eliminated from consideration for 7 exoplanet hosts, while all spectral types could be discounted for 7 host stars. The remaining 6 host stars had some but not all of the various secondary types ruled out. Because of the small sample size, we did not expect to find any stellar companions masquerading as planets, as the probability of a moderate- to high-inclination planet mimicking a face-on stellar companion is very low. Our contribution was eliminate the possibility of certain secondary spectral types for the host stars."
    },
    "1005/1005.5384_arXiv.txt": {
        "abstract": "We present a new very fast tree-code which runs on massively parallel Graphical Processing Units (GPU) with NVIDIA CUDA architecture. The tree-construction and calculation of multipole moments is carried out on the host CPU, while the force calculation which consists of tree walks and evaluation of interaction list is carried out on the GPU. In this way we achieve a sustained performance of about 100GFLOP/s and data transfer rates of about 50GB/s. It takes about a second to compute forces on a million particles with an opening angle of $\\theta \\approx 0.5$. The code has a convenient user interface and is freely available for use\\footnote{{\\tt http://castle.strw.leidenuniv.nl/software/octgrav.html}}. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.1044_arXiv.txt": {
        "abstract": "Motivated by possible astrophysical and biological applications we calculate visible and near UV spectral lines of proflavine (C$_{13}$H$_{11}$N$_3$, 3,6-diaminoacridine) in vacuum, as well as its anion, cation, and dication. The pseudopotential density functional and time-dependent density functional methods are used. We find a good agreement in spectral line positions calculated by two real-time propagation methods and the Lanczos chain method. Spectra of proflavine and its ions show characteristic UV lines which are good candidates for a detection of these molecules in interstellar space and various biological processes. ",
        "introduction": "\\label{intro} More than 150 molecules have already been detected in interstellar space \\cite{Snow}. The number of atoms in by now discovered interstellar molecules is in the range from two to more than ten. Broad absorption spectral structures have been observed in the near-UV, visible, and near-IR spectra of some stars. They are called diffuse interstellar bands (DIBs) \\cite{Snow,Herbig,Sarre,IAU}. Up to now, around 300 DIBs have been measured. Although the first such lines were detected almost ninety years ago, their carriers are still unknown. Many DIBs carriers have been suggested, but two main groups are processes on the astrophysical dust grains and molecular transitions. The molecular origin of DIBs is more dominant hypothesis in the recent literature.  Various simple or complex carbon containing molecules have been suggested as possible DIBs carriers. For example, small linear carbon chains, fullerenes, and especially polycyclic aromatic hydrocarbons (PAHs)  have been investigated as candidates for DIBs carriers \\cite{IAU,Tielens,Malloci,Mulas,Weisman}. Although various neutral and ionized PAHs, small and large ones,  have been extensively studied, the analysis still did not unambiguously confirm agreement between any specific PAHs spectra and DIBs. A few narrow DIBs have been observed in the Red Rectangle nebula \\cite{Miles}. Substructures in prominent bands of some narrow DIBs have been seen \\cite{Kerr,Ehrenfreund}. Features in these resolved profiles are consistent with rovibrational transitions in organic molecules. It is also important to study various molecules  where atoms of other chemical elements substitute C and H atoms in PAHs. Nitrogen is a very important component of biological materials, and it is one of the most abundant chemical elements in the Universe. Several molecules that contain nitrogen have already been detected in  meteorites, interstellar dust particles, and Space, or they have been analyzed by computational methods \\cite{Clemett,Bauschlicher,Hudgins}.    Proflavine, C$_{13}$H$_{11}$N$_3$  (3,6-diaminoacridine), is a molecule well known in molecular biology \\cite{Lerman,Schulman,Jain}. This molecule interacts with nucleic acids, intercalates into DNA, and blocks activities of many  viruses. C$_{13}$H$_{11}$N$_3$  is a typical example of amino-acridines which have been tested and used as anticancer drugs. It has been suggested that proflavine is one of molecules which may act as ``midwifes'' \\cite{Jain}. It is possible that this type of molecules increased a probability of DNA and RNA synthesis when biological matter was formed. For example, proflavine increases thousand times a kinetics of the synthesis of two small DNA or RNA molecules into larger one. Therefore, proflavine is one of the candidates for molecules which induced the origin of life on Earth \\cite{Thaddeus}. Optical spectra present an efficient way to detect proflavine, its ions, and derived nanostructures in biological materials and processes. In a comparison with PAHs (for example oligoacenes), available experimental and theoretical data for proflavine are scarce. In ref. 3 it has been suggested that dyes are good candidates for studies of a DIBs problem and, as a support of this idea, it was stated that proflavine absorbs at 445 nm in water at pH 7. However, existing measurements of proflavine spectra are done in aqueous and other solutions \\cite{Schulman,Mataga,Neumann,deSilvestri,Pileni}. It has been found that the optical spectrum of proflavine is very sensitive to the type of a solvent and pH factor \\cite{Mataga,deSilvestri}. There is a possibility that spectra of proflavine and its ions in vacuum and under astrophysical conditions are substantially different from ones in solutions. In addition, it is known that many materials from  the group of organic dyes show strong lines in a visible spectrum. For example, absorption lines of basic dyes such as methylene blue, thionine, neutral red, safranine, acridine orange, acridine yellow, and proflavine are all in an interval from 400 \\AA\\, to 800 \\AA\\, \\cite{Neumann}. Although many organic dyes are good candidates for DIBs studies, because of a ``midwife'' idea \\cite{Jain} and possible applications in astrobiology, we have selected proflavine for our study.    It is known that many molecules in interstellar space exist in its ionized state. It has been found that positive and negative ions of PAH molecules play an important role in the chemistry of diffuse interstellar clouds \\cite{Lepp}. Therefore, it is also important to analyze ions of proflavine.  We use pseudopotential density functional (DFT) methods to compute the ground-state properties of proflavine and its ions, as well as the pseudopotential time-dependent density functional theory (TDDFT) to calculate their optical spectra. Pseudopotential methods are computationally less demanding than all-electrons methods and offer a possibility to study larger systems. Methods based on the pseudopotential DFT are now substantially developed to provide reliable results for structural, electronic and optical properties of various materials \\cite{Martin}. TDDFT \\cite{Runge} is recently used to describe optical properties of many materials \\cite{Yabana,Malloci,Mulas,Weisman,Lopez,Varsano,Rocca}. It is known that electronic spectra of PAHs are in the UV and visible part of the spectrum. Optical spectra of PAHs and many other organic molecules in this region of the spectrum are well described by TDDFT \\cite{Malloci,Mulas,Weisman,Lopez,Varsano,Rocca}. ",
        "conclusions": "\\label{concl}  We calculated optical spectra of proflavine, its anion, cation, and dication using pseudopotential density functional theory methods. A comparison of three TDDFT pseudopotential methods shows that they produce spectra of proflavine which  qualitatively agree well, and line positions differ 0.14 eV, at the most. The difference of an average result (calculated by three TDDFT methods) for a visible line position and a corresponding average experimental value for a neutral proflavine molecule in water \\cite{deSilvestri,Pileni} is also 0.14 eV. The shapes of spectra for proflavine and its ions are rather similar to each other, but the change in the line positions induced by a different charge is apparent. UV lines are characteristic for spectra of proflavine and its ions and give an efficient way for their detection, either in Space, meteorites, astrophysical dust particles, or in biological processes. However, the UV spectral region is much less studied than visible one. It is important to analyze available data and to perform new studies of substituted polycyclic aromatic hydrocarbons. We hope that our work will generate an interest for measurements of  optical spectra of proflavine and other organic dyes in the gaseous phase and under various astrophysical conditions."
    },
    "1005/1005.3041_arXiv.txt": {
        "abstract": "Recent numerical and analytical studies have shown that galaxies accrete most of their baryons via the cold mode, from streams with temperatures $T\\sim10^{4}-10^{5}$ K. At these temperatures, the streams should radiate primarily in the \\Lya~line and have therefore been proposed as a model to power the extended, high-redshift objects known as \\Lya~blobs, and may also be relevant for powering a range of less luminous \\Lya~sources. We introduce a new \\Lya~radiative transfer code, $\\alpha RT$, and calculate the transport of the \\Lya~emission from cold accretion in cosmological hydrodynamical simulations. In this paper, we describe our methodology, and address physical and numerical issues that are critical to making accurate predictions for the cooling luminosity, but that have been mostly neglected or treated simplistically so far. In particular, we highlight the importance of self-shielding and of properly treating sub-resolution models in numerical simulations. Most existing simulations do not self-consistently incorporate these effects, which can lead to order-of-magnitude errors in the predicted cooling luminosity. Using a combination of post-processing ionizing radiative transfer and re-simulation techniques, we develop an approximation to the consistent evolution of the self-shielded gas. We quantify the dependence of the \\Lya~cooling luminosity on halo mass at $z=3$ for the simplified problem of pure gas accretion embedded in the cosmic radiation background and without feedback, and present radiative transfer results for a particular system. While pure cooling in massive halos (without additional energy input from star formation and AGN) is in principle sufficient to produce $L_{\\alpha}\\sim10^{43}-10^{44}$ erg s$^{-1}$ blobs, this requires including energy released in gas of density sufficient to form stars, but which is kept 100\\% gaseous in our optimistic estimates. Excluding emission from such dense gas yields lower luminosities by up to one to two orders of magnitude at high masses, making it difficult to explain the observed \\Lya~blobs with pure cooling. Resonant scattering produces diffuse \\Lya~halos, even for centrally concentrated emission, and broad double peaked line profiles. In particular, the emergent line widths are in general not representative of the velocity dispersion within galactic halos and cannot be directly used to infer host halo masses. ",
        "introduction": "\\label{introduction} \\cite{2000ApJ...532..170S} discovered two extremely bright, large, and diffuse \\Lya-emitting `blobs' in a narrowband survey of the SSA 22 proto-cluster region at $\\langle z \\rangle=3.09$ \\citep[for the early detection of extended \\Lya~emission at $z\\sim2.4$, see also][]{1996ApJ...457..490F, 2001ApJ...554.1001F, 1999AJ....118.2547K}. These two blobs, labelled LAB1 and LAB2, have physical extent $\\gtrsim140$ kpc, are more luminous ($L_{\\rm Ly\\alpha}\\sim10^{44}$ erg s$^{-1}$) than typical line emitters at the same redshift by a factor $\\sim10-100$, and unlike similar halos around radio galaxies have no detectable radio continuum. Since their discovery, these blobs have become some of the most spectacular of a new class of sources that now counts several tens of members \\citep[e.g.,][]{2004AJ....128..569M, 2004ApJ...602..545P, 2004ApJ...614...75F, 2005ApJ...629..654D, 2006A&A...452L..23N, 2006ApJ...648...54S, 2007MNRAS.378L..49S, 2008ApJ...678L..77P, 2009ApJ...702..554P, 2009ApJ...696.1164O, 2009ApJ...693.1579Y, 2010ApJ...719.1654Y} and whose nature remains unclear. Because detecting spatially extended \\Lya~emission usually requires deep narrowband imaging, which covers thin redshift slices, only a small volume the Universe has been effectively surveyed for them to date; the blobs and their fainter analogues are therefore likely to be numerous and cosmologically significant. In fact, existing observations and theoretical models indicate that they may be the sites of massive galaxy formation and also may display signatures of associated active galactic nuclei (AGN) and/or supernova feedback. The \\Lya~blobs thus provide a unique opportunity to probe the processes driving and regulating galaxy formation, and may be intimately related to phenomena including proto-clusters, mergers, and submillimeter galaxies.\\\\ \\\\ Although the most extreme \\Lya~blobs have received the most attention, there in fact exists a wide continuum of spatially extended \\Lya~sources at high redshift, for example the ones discovered by \\cite{2006ApJ...648...54S} with line luminosities $\\sim10^{42}$ erg s$^{-1}$ and those discovered by \\cite{2008ApJ...681..856R}, with line luminosities as low as $\\sim10^{39}$ erg s$^{-1}$. Understanding the nature of these fainter but more numerous sources is equally important to develop a physical picture of galaxy formation. In fact, the fainter sources likely probe different (perhaps earlier) stages of galaxy assembly.\\\\ \\\\ A central puzzle for the \\Lya~blobs is that many of them do not appear to have a central source energetic enough to power their entire \\Lya~emission \\citep[e.g.,][]{2004AJ....128..569M, 2006A&A...452L..23N}. Even for the blobs which do have energetically-sufficient counterparts \\citep[e.g.,][]{2005MNRAS.363.1398G, 2007ApJ...655L...9G, 2009ApJ...700....1G, 2009ApJ...692.1561W}, it is unclear whether that energy can actually couple effectively to the \\Lya~emission. Submillimeter starbursts and obscured AGN imply the presence of large quantities of dust, which acts to destroy \\Lya~photons particularly efficiently \\citep[][]{1990ApJ...350..216N}. It is therefore uncertain whether \\Lya~photons produced by such dust-enshrouded sources can escape in significant amounts. Nevertheless, several different mechanisms been proposed to power the blobs and can be broadly divided into three categories:\\\\ \\\\ {\\bf Embedded star formation or AGN}, possibly obscured from direct view by dust, could photoionize the surrounding hydrogen nebula \\citep[][]{1998MNRAS.299..661M, 2001ApJ...556...87H, 2004Natur.430..999W, 2005A&A...436..825W, 2007ApJ...657L..69L}.\\\\ \\\\ {\\bf Superwinds} driven by starburst supernovae could explain the observed sizes and kinematics of the blobs, with the \\Lya~emission being generated in the swept up material \\citep[][]{2000ApJ...532L..13T, 2001ApJ...562L..15T, 2003ApJ...591L...9O, 2004ApJ...613L..97M, 2005Natur.436..227W, 2005MNRAS.363.1398G}. AGNs could also drive similar winds \\citep[e.g.,][]{2005ApJ...618..569M, 2005MNRAS.361..776S, 2005Natur.433..604D, 2006ApJS..166....1H}.\\\\ \\\\ {\\bf Cooling radiation}, emitted as gas accretes onto forming galaxies, could produce luminous and extended structures. If a large fraction of the accreting gas has a temperature $T\\sim10^{4}-10^{5}$ K, then most of the cooling radiation could be \\Lya~\\citep[][]{1977MNRAS.180..479F, 1991ApJ...377..365K, 1992ApJ...391..608H, 2000ApJ...537L...5H, 2001ApJ...562..605F, 2003MNRAS.345..349B, 2005MNRAS.363....2K, 2005ApJ...622....7F, 2006ApJ...649...14D, 2006ApJ...640..539Y, 2009MNRAS.400.1109D, 2010MNRAS.407..613G, 2010MNRAS.402.1449D}.\\\\ \\\\ In this work, motivated by recent progress on our understanding of how galaxies get their gas, we focus on the third possibility. The methods developed could however be applied to the first two classes of models as well, and we plan to extend our calculations to model those processes in the future.\\\\ \\\\ In the classic sketch of galaxy formation \\citep[][]{1977MNRAS.179..541R, 1977ApJ...211..638S, 1978MNRAS.183..341W}, gas falling into dark matter halos is shocked and heated to the virial temperature. For a galaxy with a mass similar to that of the Milky Way, the shocked gas attains a temperature $T_{\\rm vir}\\sim10^{6}$ K. In the dense inner regions of the halos, this gas efficiently radiates its thermal energy, loses its pressure support, and settles into compact discs where it can form stars. A wealth of recent work however suggests that this picture requires an important modification: most of the gas is never strongly shocked as it flows toward the central forming galaxy, but rather accretes in a ``cold mode'', maintaining a temperature $T\\lesssim10^{5}$ K. Moreover, this cold accretion proceeds through dense filaments rather than in a spherically symmetric fashion.\\\\ \\\\ Although the importance of cold accretion in galaxy formation has only recently been demonstrated in high-resolution three-dimensional hydrodynamical simulations \\citep[e.g.,][]{2001ApJ...562..605F, 2003ASSL..281..185K, 2005MNRAS.363....2K, 2009MNRAS.395..160K, 2008MNRAS.390.1326O, 2009ApJ...694..396B, 2009Natur.457..451D}, \\cite{1977ApJ...215..483B} had argued on the basis of analytic models of protogalaxy collapse that the amount of shock heating could be small for plausible physical conditions. Moreover, already in the first simulations of forming galaxies, most of the gas never heated above $T\\sim3\\times10^{4}$ K \\citep[][]{1991ApJ...377..365K} and the importance of filamentary structures was recognized by \\cite{1993ApJ...412..455K} and \\cite{1994MNRAS.270L..71K}. \\cite{2003MNRAS.345..349B} carried out a stability analysis, supported by one-dimensional hydrodynamical simulations \\citep[see also][]{2007MNRAS.380..339B, 2006MNRAS.368....2D}, and found that when the radiative cooling is efficient compared with the infall rate, the post-shock gas becomes unstable and cannot support the shock. When applied to cosmology, their results agree well with those of three-dimensional simulations.\\\\ \\\\ The \\Lya~emission from cold accretion has already been the subject of some studies. \\cite{2001ApJ...562..605F} first evaluated the \\Lya~cooling luminosity from hydrodynamical simulations and suggested that it could account for the \\Lya~blobs discovered by \\cite{2000ApJ...532..170S}. \\cite{2000ApJ...537L...5H} reached a similar conclusion using simplified analytic arguments. Also using simulations, \\cite{2005ApJ...622....7F} found that the \\Lya~cooling radiation from structure formation could account for some, but not all, of the luminosity of the \\Lya~blobs. \\cite{2006ApJ...640..539Y} studied both the hydrogen and helium cooling radiation using simulations, but found the hydrogen \\Lya~luminosity to be strongly dependent on the self-shielding correction applied. Recently, \\cite{2009MNRAS.400.1109D} developed an analytic model and suggested that cooling radiation from the cold mode could account for all the \\Lya~blobs under reasonable assumptions. Using adaptive mesh refinement (AMR) simulations, \\cite{2010MNRAS.407..613G} provided supporting evidence for this picture. In our discussion (\\S \\ref{discussion}), we will contrast our main results with those of \\cite{2010MNRAS.407..613G}, concluding that the differences with theirs most likely originate from the treatments of self-shielding and sub-resolution modeling, which are a focus of our study.\\\\ \\\\ No study focusing specifically on cooling emission has however combined realistic \\Lya~radiative transfer with hydrodynamical simulations before.\\footnote{Other authors have included a cooling component in radiative transfer calculations of \\Lya-emitting galaxies \\citep[e.g.,][]{2006ApJ...645..792T, 2007ApJ...657L..69L, 2009ApJ...696..853L, 2009ApJ...696..853L}, but have not explicitly separated out the signatures of pure cooling or investigated the important uncertainties in detail.}Because \\Lya~photons resonantly scatter, the resultant morphologies, spatial extents, and spectra are strongly modified by radiative transfer effects \\citep[e.g.,][]{2006ApJ...649...14D}. Since the \\Lya~photons tend to follow paths of least resistance in space and frequency (\\S \\ref{lya radiative transfer}), the spatial geometry and bulk velocity fields play critical roles in determining the radiation transport \\cite[for observational evidence of these effects, see e.g.][]{1998A&A...334...11K, 2003ApJ...598..858M}. Fully three-dimensional calculations are therefore necessary to make realistic predictions. Moreover, both the existing analytical and numerical studies have limitations that could induce important errors in quantities as basic as the integrated \\Lya~luminosity of the cold streams. One such uncertainty arises from the exponential dependence of the \\Lya~emissivity on the gas temperature for the temperatures $T \\sim 10^{4}$ K that are characteristic of cold accretion (\\S \\ref{lya emission}). At present, most galaxy formation simulations do not self-consistently predict the temperature distribution within the streams. In fact, existing simulations usually do not follow the transport of the ultra-violet (UV) radiation that ionizes and heats dense gas. This seriously limits the predictive power of these calculations, since small errors in the temperatures can result in large errors in the \\Lya~cooling emission, and potentially grossly violate energy conservation. As we will demonstrate, models of sub-resolution physics in hydrodynamical simulations can also introduce large errors if not properly taken into account. Analytical studies based on energetic considerations are not as sensitive to the temperature of the cold streams \\citep[e.g.,][]{2009MNRAS.400.1109D}, but are not immune of uncertainties either, since they rely on assumptions regarding the efficiency of \\Lya~emission. Moreover, their simplified nature does not lend itself to detailed radiative transfer predictions. Resolving these issues is critical to relating the \\Lya~emission from cold accretion to observations.\\\\ \\\\ Our ultimate goal is a systematic investigation of the \\Lya~emission from galaxy formation that is both detailed in its predictions, and robust. By detailed, we envision predictions that can be directly compared with observations, and therefore require both 3D hydrodynamical simulations and realistic radiative transfer. By robust, we mean that the predictions should be free of assumptions that introduce the kind of large uncertainties that existing studies are subject to. Due to the complexity of the problem, this ultimate goal is likely to require a long-term effort. The present paper is dedicated to laying down some of the foundations for this research program. Specifically, we present a new \\Lya~radiative transfer code, named $\\alpha RT$, and describe its application to the cooling radiation in cosmological simulations of galaxy formation \\citep[for other applications of \\Lya~radiative transfer codes to hydrodynamical simulations, see e.g.][]{2005ApJ...628...61C, 2006ApJ...645..792T, 2007ApJ...657L..69L, 2009ApJ...696..853L, 2010ApJ...708.1048K, 2010ApJ...716..574Z}. We pay particular attention to clarifying the physical and numerical uncertainties of these calculations, in particular with respect to the predicted \\Lya~luminosities, and illustrate the importance of radiative transfer effects. To do so, aside for calculating the \\Lya~luminosities of a sample of halos from a cosmological volume, we focus our radiative transfer calculations on a particular system at $z=3$ and explore variations in both the emission and radiative transfer physics. We limit ourselves to the most basic physical problem of accreting halos embedded in a cosmic ionizing background and neglect feedback processes. Follow up studies will build on the results obtained here and investigate the properties of the \\Lya~emission as a function of halo mass and redshift, and will extend them by incorporating additional physics, including feedback \\citep[][]{cfcosmo}.\\\\ \\\\ We begin by describing our hydrodynamical simulations in \\S \\ref{simulations}. We address the emission of \\Lya~photons in \\S \\ref{lya emission} and explicitly demonstrate the sensitivity of the predicted \\Lya~cooling luminosity on assumptions regarding the thermal state of the self-shielded gas. Using a combination of post-processing ionizing radiative transfer and re-simulation techniques, we develop an approximation to consistently model the evolution of the dense gas in the hydrodynamical simulations, resulting in the most robust numerical predictions to date. In \\S \\ref{lya radiative transfer}, we present the results of \\Lya~radiative transfer calculations for a particular system of total mass $M_{h}=2.5\\times10^{11}$ M$_{\\odot}$ at $z=3$ and highlight the role of both the bulk velocity flows and of the resonant scatters in shaping the emergent morphology and spectrum. Finally, we discuss our results and conclude in \\S \\ref{discussion}. The Appendices document the radiative transfer code $\\alpha RT$ introduced in this work, the ionizing radiative transfer method, and relevant analytical estimates.\\\\ \\\\ Throughout, we assume a cosmology with $(\\Omega_{\\rm m},~\\Omega_{\\rm b},~\\Omega_{\\Lambda},~h,~\\sigma_{8},~n_{\\rm s})=(0.28,~0.046,~0.72,~0.70,~0.82,~0.96)$, as inferred from the \\emph{Wilkinson Microwave Anisotropy Probe} (WMAP) five-year data in combination with baryon acoustic oscillations and supernovae \\citep[][]{2009ApJS..180..330K}. While some of our hydrodynamical simulations were run with slightly different parameters, none of our conclusions are sensitive to the details of the cosmology. We assume hydrogen and helium mass fractions of $X=0.75$ and $Y=0.25$ \\citep[e.g.,][]{2001ApJ...552L...1B}, the collisional ionization coefficients given in \\cite{1996ApJS..105...19K}, the \\Lya~collisional excitation coefficient and average number of \\Lya~photons produced per recombination from \\cite{2006agna.book.....O}, and the recombination coefficients in the appendix of \\cite{1997MNRAS.292...27H}. For convenience, some symbols used in this work are defined in Table \\ref{definitions}.  \\begin{deluxetable*}{ll} \\tablewidth{0pc} \\tablecaption{Symbols used in this work\\label{definitions}} \\tabletypesize{\\footnotesize} \\tablehead{\\colhead{Symbol} & \\colhead{Definition}} \\startdata $n_{i}$                   & number density of species $i$ \\\\ $N_{i}$                   & column density of species $i$ \\\\ $T$                       & gas temperature \\\\ $\\tau_{\\nu}$              & optical depth at frequency $\\nu$ \\\\ $\\nu_{0}$                 & \\Lya~central frequency \\\\ $\\Delta \\nu_{\\rm D}$      & \\Lya~Doppler width \\\\ $x$                       & dimensionless frequency offset $(\\nu-\\nu_{0})/\\Delta \\nu_{\\rm D}$ \\\\ $\\Gamma_{i}$              & photoionization rate of species $i$ \\\\ $\\Gamma_{i,c}$            & collisional ionization coefficient of species $i$ \\\\ $C_{\\rm Ly\\alpha}$         & \\Lya~collisional excitation coefficient \\\\ $\\epsilon_{\\alpha}$       & \\Lya~emissivity \\\\ $\\alpha_{i}^{A,B}$\\tablenotemark{a}        & case A, B recombination coefficients to species $i$ \\\\ $x_{\\rm HI}, x_{\\rm HII}$              & fractions of hydrogen in HI, HII \\\\ $y_{\\rm HeI},~y_{\\rm HeII},~y_{\\rm HeIII}$ & fractions of helium in HeI, HeII, HeIII \\enddata \\end{deluxetable*} ",
        "conclusions": "\\label{discussion} Motivated by observations of a variety of extended \\Lya~sources at high redshift and by recent studies suggesting that cold mode accretion could account for the majority of them \\citep[e.g.,][]{2009MNRAS.400.1109D}, we have used hydrodynamical simulations to predict the \\Lya~cooling emission of forming galaxies. We have introduced a new \\Lya~radiative transfer code, $\\alpha RT$, and for the first time applied such a code to the particular problem of cooling radiation in cosmological hydrodynamical simulations. In this work, which will serve as the foundation for follow up studies of the astrophysical implications, we have quantified the theoretical and numerical uncertainties in predicting the properties of \\Lya~cooling radiation. To do so, we have considered the simplest physical problem of galaxy assembly embedded in the cosmic UV background but without feedback processes. We have shown that the \\Lya~cooling luminosities, morphologies, and spectra of the cold streams are strongly dependent on the treatment of the ionization and thermal state of the emitting gas, and that naive assumptions can produce artificially high cooling luminosities. Previous studies have usually relied on such assumptions, because most existing hydrodynamical simulations do not self-consistently follow the transfer of ionizing radiation \\citep[but see][]{2006ApJ...644L...1S, 2007ApJ...657L..69L, 2009ApJ...704.1640L}, whereas the cold streams in reality self-shield. The predicted cooling luminosity is so sensitive to the thermal state of the gas principally because of the exponential dependence of the collisional excitation coefficient on temperature in the range $T\\sim10^{4}-10^{5}$ K characteristic of the cold streams. We have also demonstrated that subresolution physics models, when not carefully taken into account, can lead to large errors as well.\\\\ \\\\ Having explicitly illustrated the range of uncertainties, we have made a systematic attempt to converge on the correct answer. Our strategy consisted of first post-processing our hydrodynamical simulations with ionizing radiative transfer to identify the self-shielded regions. By inspection, we found that there is a fairly well defined total hydrogen density $n_{\\rm H}\\sim0.01$ cm$^{-3}$ above which the gas self-shields, at least at the redshifts $z \\sim 3$ characteristic of the existing observations. We have then rerun simulations with exactly the same initial conditions and other physical parameters, but with the ionizing background turned off in regions exceeding the fiducial density threshold $n_{\\rm H}=0.01$ cm$^{-3}$. By turning the ionizing background off in dense regions on the fly, this dense gas is consistently evolved with modified heating and cooling functions, as well as the corresponding dynamical response. These simulations provide us with our best estimates for the actual \\Lya~cooling luminosity of the cold streams, in good agreement with energetic considerations.\\\\ \\\\ For our \\Lya~radiative transfer calculations, we focused on a particular halo (A1) of total mass $M_{h}=2.5\\times10^{11}$ M$_{\\odot}$ at $z=3$. We have shown that it is not only the integrated \\Lya~luminosity of the system that depends strongly on assumptions regarding the thermal state of the gas, but also its apparent morphology and spectrum. To isolate the role of different physical effects in shaping the observational properties, we have sequentially turned off separate pieces of physics. Both the \\Lya~resonant scatters as well as the bulk velocity structure of the system are critical in shaping the emergent spectrum. The resultant effective line width in general differs strongly from the optically thin expectation owing to these radiative transfer effects, and we therefore caution against measuring the masses of systems based on the \\Lya~line alone. The intrinsically extended nature of the cooling emission and the spatial diffusion owing to resonant scattering combine to produce objects that are spatially extended on the sky, but with an apparent size that depends sensitively on the observational surface brightness threshold. As a result, only the central few kpc of the A1 halo \\Lya~cooling emission at $z=3$ sticks out above the surface brightness threshold $\\sim10^{-18}$ erg s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$ typically achieved to date (Fig. \\ref{A1 RT vs physics conditional}). If the characteristic size of the \\Lya~emission scales as $R\\propto M_{\\rm h}^{1/3}$, then the surface brightness should scale as $SB \\propto L_{\\alpha} / R^{2} \\propto L_{\\alpha} / M_{\\rm h}^{2/3}$. For our consistent prescriptions 7$-$9, Figure \\ref{Lalpha vs M} shows that the high-mass luminosity slope varies from about 3/4 to 1 in the most optimistic case, i.e. $SB \\propto M_{\\rm h}^{-0.33...-0.08}$, so that even the higher-mass systems are unlikely to show up as sources with extents $\\sim100$ kpc in existing observations from cooling luminosity alone. Fainter sources may however well be detectable in deeper observations, such as the 100-hr long-slit spectrum reported by \\cite{2008ApJ...681..856R}.\\\\ \\\\ Our analysis of a sample of halos from a cosmological volume at $z\\approx3$ confirms that \\citep[in agreement with previous studies;][]{2000ApJ...537L...5H, 2001ApJ...562..605F, 2009MNRAS.400.1109D, 2010MNRAS.407..613G} cooling emission alone can \\emph{in principle} produce luminosities $L_{\\alpha} \\sim 10^{43}-10^{44}$ erg s$^{-1}$ sufficient to account for luminous \\Lya~blobs. This however requires quite optimistic assumptions under which most of the energy is released close to, or within, the accreting discs, at densities sufficient to form stars according to the observed Kennicutt-Schmidt relation. In those optimistic predictions based on prescription 7, star formation was artificially turned off (yielding purely gaseous discs), so that it is unlikely that this entire cooling emission can be realized in reality without being overwhelmed by stellar emission.\\footnote{In their simulations of LBGs at $z=3.6$, Laursen et al. 2009a find that only $\\sim10$\\% of the total \\Lya~emission comes from cooling, although we show in Appendix \\ref{analytic considerations appendix} that this ratio should depend on halo mass and redshift.} Our findings in this respect are at odds with the recent simulations of \\cite{2010MNRAS.407..613G}, who argue that pure cooling radiation can explain the observed giant \\Lya~blobs. Our investigation of the sources of error in the \\Lya~luminosity predicted from simulations (\\S \\ref{lya emission}) allows us to understand the differences. First, we have demonstrated that the predicted \\Lya~luminosity is very sensitive to the treatment of self-shielded gas and that at $z\\approx3$, the characteristic density above which gas self-shields from the UVB is $n_{\\rm H}\\approx0.01$ cm$^{-3}$. \\cite{2010MNRAS.407..613G} did not explicitly perform ionizing radiative transfer and instead assumed a higher self-shielding threshold of $n_{\\rm H}=0.1$ cm$^{-3}$. This higher density threshold translates into an overestimate of the \\Lya~luminosity owing to artificial photoheating of gas in the density range $0.01\\leq n_{\\rm H} \\leq 0.1$ cm$^{-3}$. Examination of the luminosity-weighted $T-n_{\\rm H}$ histogram in their Figure 7 in fact indicates that the bulk of the \\Lya~luminosity they predict originates from this density regime (including in clumps that would correspond to satellite galaxies, even when the central galaxy is excised), whereas \\Lya~emission is strongly suppressed at these densities in our approximation to the consistent thermal evolution of self-shielded gas, owing to rapid cooling below $10^{4}$ K (Fig. \\ref{T vs nH}). While there is some uncertainty in the precise value of the self-shielding threshold, our radiative transfer calculations show that it is clearly below 0.1 cm$^{-3}$ at the redshifts under consideration (Fig. \\ref{xHI vs nH}), and the value $n_{\\rm H}\\approx0.01$ cm$^{-3}$ is further supported by recent radiative transfer calculations by other groups \\citep[][]{2010arXiv1006.5345N, 2010arXiv1004.2503A}. The cleanest contrast between our results and those of \\cite{2010MNRAS.407..613G} is provided by their predictions of the \\Lya~luminosity originating from the ``streams'' alone, in which they excluded all emission from within 20\\% of the virial radius. Those can be compared, at best, with our pessimistic prescriptions 8 and 9, in which only gas with $n_{\\rm H}>0.13$ cm$^{-3}$ (above the subresolution star formation threshold) is excluded; our predictions are in those cases a factor $\\gtrsim10$ lower than theirs. It would not be fair to compare our most optimistic prescription 7 with the \\cite{2010MNRAS.407..613G} results with the central 20\\% of the virial radius excluded, since most of the extra emission under this prescription occurs very close to, or inside, the central galaxy. Finally, the simulations of \\cite{2010MNRAS.407..613G} included supernova feedback, which provides an additional source of energy and could also contribute significantly to boosting the cooling luminosity, even in the absence of ionizing photons from associated local sources  \\citep[e.g.,][]{2000ApJ...532L..13T, 2001ApJ...562L..15T, 2003ApJ...591L...9O, 2004ApJ...613L..97M}.\\\\ \\\\ It is further notable that many \\Lya~blobs show spectral signatures of outflows (dominant red peak), at odds with the expectation of a more prominent blue peak in the case of pure cooling of infalling gas.\\footnote{There are exceptions, for example certain regions of the LAB2 blob at $z=3.09$ \\citep[][]{2005Natur.436..227W}.} Nonetheless, the release of gravitational potential energy through cooling radiation at observationally interesting levels is an inevitable prediction of galaxy formation in $\\Lambda$CDM and the growing body of theoretical work supporting the cold mode scenario implies that a significant fraction would come out in \\Lya. Such sources are therefore poised to be routinely detected in future surveys, and likely account for at least a subpopulation of the more modest extended \\Lya~sources, such as the ones detected by \\cite{2008ApJ...681..856R}. In fact, at least three of the \\cite{2008ApJ...681..856R} sources (\\#15, 36, and 37) show spectral signatures suggestive of infall (with a more prominent blue peak). Interestingly, GALEX non-detections indicate that the bright \\Lya~blobs are much rarer at $z=0.8$ than at $z \\gtrsim 2$ \\citep[][]{2009AJ....138..986K}. This redshift evolution is reminiscent of the gradual disappearance of the cold mode at low redshift predicted by theoretical studies \\citep[e.g.,][]{2003MNRAS.345..349B, 2006MNRAS.368....2D, 2005MNRAS.363....2K, 2009MNRAS.395..160K}. It is important to note that the cold mode could play an important role in the existence and properties of the \\Lya~blobs even if they are not energetically dominated by cooling emission. In fact, the presence of cold neutral gas in galactic halos enhances the conversion stellar or AGN power into \\Lya~photons, and scattering off such gas may be necessary to produce the morphological and spectral properties observed.\\\\ \\\\ Although we have made important progress in accurately predicting the \\Lya~properties of cold accretion, much work remains to be done. For instance, we have focused on the simplest physical problem of accretion embedded in the cosmic UV background and neglected feedback processes. In the actual Universe, galactic winds are observed to be ubiquitous at high redshift \\citep[e.g.,][]{2003ApJ...588...65S, 2010arXiv1003.0679S} and simulations suggest that they are in fact needed to reproduce the observed baryonic mass function of galaxies \\citep[e.g.,][]{2009MNRAS.396.2332K, 2010MNRAS.406.2325O}. Furthermore, AGN can inject large amounts of energy in the surrounding gas \\citep[][]{2005MNRAS.361..776S, 2006ApJS..166....1H, 2009arXiv0909.2872D}, which could radiate even after the nucleus has shut off. Such processes will modify the kinematic and thermal properties of the circumgalactic medium, and therefore its \\Lya~signatures. Moreover, they provide additional sources of energy that could enhance the total emission. The ionizing radiation produced by embedded star formation or AGN \\citep[e.g.,][]{2010ApJ...708.1048K}, the effects of metals on the cooling function \\citep[e.g.,][]{2010MNRAS.403L..16C}, and the destruction of \\Lya~photons by dust \\citep[as shown by][]{2009ApJ...704.1640L} are also likely to be important to reproduce the observed sources, but have not been modeled here for simplicity. Since cold accretion likely fuels star formation in halo centers, at rates similar to the cold gas accretion rates \\citep[e.g.,][]{2005MNRAS.363....2K, 2009MNRAS.395..160K, baryona}, gas accretion and star formation are intrinsically linked and their roles in \\Lya~emission are difficult to decouple. Follow up studies will build upon the technical foundation established in this work and include some of these effects. Ultimately, cosmological statistics like the luminosity and correlation functions of the \\Lya~sources will be predicted. It will also be interesting to consider complementary observational probes of the cold streams, including their \\Lya~polarization \\citep[e.g.,][]{1999ApJ...520L..79R, 2008MNRAS.386..492D}, their absorption signatures, and other emission lines (including from metals).\\\\ \\\\ On a more basic level, our understanding of the physics of the cold mode is incomplete. In particular, it is not exactly clear how the cold streams release their gravitational energy. At least some simulations indicate that they maintain a roughly constant velocity, rather than accelerate, as they fall into galactic halos \\citep[][]{2005MNRAS.363....2K}. This implies that the streams would continuously radiate the gravitational work done on them, perhaps by undergoing a series of weak shocks. However, these weak shocks have not been explicitly identified in existing simulations. \\cite{2003MNRAS.345..349B} instead envision a scenario in which the cold streams free fall into the halos and release little energy before hitting the central disc in a strong shock, which may be partially supported by our results which suggest that a large fraction of the energy could be released near the halo center. Infalling cold gas might also release some energy through interactions with halo sub-structure and with the lower density, hot halo gas. These interactions could involve hydrodynamic instabilities below the resolution of current generation cosmological hydrodynamic simulations, both SPH and AMR. While we have taken the pragmatic point of view of predicting the observational signatures implied by our current galaxy formation models, it is conceivable that the existing simulations are subject to numerical limitations. Work is underway to elucidate some of these hydrodynamical issues using the new shock-capturing, moving mesh code AREPO \\citep[][]{2010MNRAS.401..791S}. When the capability becomes available, it would also be useful to perform truly self-consistent ionizing radiative transfer on the fly in order to definitively capture the thermal evolution of the gas."
    },
    "1005/1005.1933_arXiv.txt": {
        "abstract": "The metallicity in galaxy clusters is expected to originate from the stars in galaxies, with a population dominated by high mass stars likely being the most important stellar component, especially in rich clusters.  We examine the relationship between the metallicity and the prominence of galaxies as measured by the star to baryon ratio, M$_*$/M$_{bary}$.  Counter to expectations, we rule out a metallicity that is proportional to M$_*$/M$_{bary}$, where the best fit has the gas phase metallicity decreasing with M$_*$/M$_{bary}$, or the metallicity of the gas plus the stars being independent of M$_*$/M$_{bary}$. This implies that the population of stars responsible for the metals is largely proportional to the total baryonic mass of the cluster, not to the galaxy mass within the cluster. If generally applicable, most of the heavy elements in the universe were not produced within galaxies. ",
        "introduction": "While nucleosynthesis is undoubtedly responsible for the production of the elements, there is less certainty regarding the stellar populations involved, as well as the time and location of elemental production.  Historically, these issues have been studied within individual galaxies, an ideal approach if they are nearly closed box systems.  However, studies find that galaxies have lost the majority of their baryons \\citep{mcga05} and some undetermined amount of metals.  For example, a galaxy like the Milky Way is missing 75\\% of its original baryon content while for M33, the deficit is about 90\\% \\citep{ande10}.  Galaxy clusters provide another useful system for studying metallicity evolution, as moderate and rich clusters retain most of their baryons (e.g., \\citealt{bohr09}).  Baryons lost from galaxies are captured in the hot cluster medium, which is effectively studied in X-rays (e.g., \\citealt{wern08}). The first order of business is to assess whether the stars in galaxies have produced the observed metallicity.  Simple estimates indicate a problem in that there are too few stars to account for the heavy element mass \\citep{arim97,port04,loew06}.  We estimate that the metallicity of stars in galaxies, averaged over a standard luminosity function, is about 0.5 solar. Galaxies comprise about 5-10\\% of the baryons in rich clusters, so if the metals in these stars were spread through the cluster, we expect the hot gas metallicity to be about 5\\% solar, but the typical metallicity is 40\\% solar (e.g., \\citealt{snow08}), about an order of magnitude higher.  This issue was examined by several authors \\citep{port04,loew06}, who did careful calculations of metallicity production from the various types of supernovae, using a cosmic star formation history (review of techniques by \\citealt{borg08}).  \\citet{loew06} concludes that with any of the standard initial mass functions, the metals produced by the stars in galaxies cannot produce the cluster metals.  He posits the existence of a population of mainly high-mass stars to produce the cluster metallicity and he investigates the time evolution. \\citet{port04}  reach a similar conclusion and adopt the IMF proposed by \\citet{arim97}, which is top-heavy (\\citealt{fabj08} use a top-heavy IMF and a normal IMF).  A dissenting view is given by \\citet{siva09} who find that for some clusters, the metals might have originated in galaxies.  However, several of their adopted values are relatively extreme compared to the work of \\citet{loew06} and \\citet{port04}.  In a related paper \\citet{ande10}, we argued that field spirals are missing their baryons due to the presence of an early stellar population dominated by high-mass stars, which heats and enriches the gas before most of it collapsed into proto-galaxies.  In this picture, which is driven by simple observational constraints, galaxies are baryon-poor because the gas never fell into galactic dark matter potential wells.  Preheating (and pre-enrichment) by this early population would need to occur near the onset of galaxy formation, which occurs before cluster formation. Therefore, this picture should apply to systems that eventually become galaxy clusters.  It is natural to expect that the mass of the high-mass weighted early population of stars is proportional to the mass of a galaxy, so the collective mass of this population in a cluster should be proportional to the total stellar mass seen in galaxies today.  Translating this into a metallicity prediction, the metallicity of the cluster should be proportional to the ratio of the stellar mass to the total baryon mass, M$_*$/M$_{bary}$. This is the same expectation if the metals were produced and expelled from the cluster galaxies.  We show that this basic expectation is excluded by the data, potentially leading to a profound change in our understanding of metallicity production in clusters and in the universe. ",
        "conclusions": ""
    },
    "1005/1005.4511_arXiv.txt": {
        "abstract": "{ B[e] supergiants are surrounded by large amounts of hydrogen neutral material, traced by the emission in the optical [O{\\sc i}] lines. This neutral material is most plausibly located within their dense, cool circumstellar disks, which are formed from the (probably non-spherically symmetric) wind material released by the star. Neither the formation mechanism nor the resulting structure and internal kinematics of these disks (or disk-like outflows) are well known. However, rapid rotation, lifting the material from the equatorial surface region, seems to play a fundamental role. } { The B[e] supergiant LHA\\,115-S\\,65 (in short: S\\,65) in the Small Magellanic Cloud is one of the two most rapidly rotating B[e] stars known. Its almost edge-on orientation allows a detailed kinematical study of its optically thin forbidden emission lines. With a focus on the rather strong [O{\\sc i}] lines, we intend to test the two plausible disk scenarios: the outflowing and the Keplerian rotating disk. } { Based on high- and low-resolution optical spectra, we investigate the density and temperature structure in those disk regions that are traced by the [O{\\sc i}] emission to constrain the disk sizes and mass fluxes needed to explain the observed [O{\\sc i}] line luminosities. In addition, we compute the emerging line profiles expected for either an outflowing disk or a Keplerian rotating disk, which can directly be compared to the observed profiles. } { Both disk scenarios deliver reasonably good fits to the line luminosities and profiles of the [O{\\sc i}] lines. Nevertheless, the Keplerian disk model seems to be the more realistic one, because it also agrees with the kinematics derived from the large number of additional lines in the spectrum. As additional support for the presence of a high-density, gaseous disk, the spectrum shows two very intense and clearly double-peaked [Ca{\\sc ii}] lines. We discuss a possible disk-formation mechanism, and similarities between S\\,65 and the group of Luminous Blue Variables. } {} ",
        "introduction": "B[e] supergiants in the Large and Small Magellanic Clouds (LMC and SMC), even though studied in great detail, are still far from being understood. A non-spherically symmetric wind is suggested by, e.g., polarimetric observations (Magalh\\~aes \\cite{Magalhaes}; Magalh\\~aes et al. \\cite{Magalhaesetal}; Melgarejo et al. \\cite{Melgarejo}), and a circumstellar disk seems to be confirmed by their strong infrared excess emission due to circumstellar dust (e.g., Zickgraf et al. \\cite{Zickgraf86}) as well as by the emission of intense molecular bands like the CO first-overtone bands in the near infrared $K$-band spectra (McGregor et al. \\cite{McGregor_a,McGregor_b, McGregor89}; Morris et al. \\cite{Morris}), and TiO bands at optical wavelengths (Zickgraf et al. \\cite{Zick89}). Based on the analysis of optical [O{\\sc i}] lines, it has recently been suggested that the disks around B[e] supergiants are neutral in hydrogen right from (or at least very close to) the stellar surface (Kraus \\& Borges Fernandes \\cite{KB}; Kraus et al. \\cite{Vlieland, KBA}), so that molecules and dust are forming in the vicinity of the star. Indeed, for the LMC B[e] supergiant R\\,126, Kastner et al. (\\cite{Kastner}) found that the inner edge of the massive dusty disk must be located at $\\sim 360\\,R_*$, which is about three times closer to the star than the value of $\\sim 1000\\,R_*$ formerly suggested by Zickgraf et al. (\\cite{Zickgraf85}). The much closer inner edge of the dust disk requests that the [O{\\sc i}] emission must originate from distances between the surface of R\\,126 and $\\sim 360\\,R_*$, which was confirmed by Kraus et al. (\\cite{KBA}) considering a hydrogen neutral disk right from the stellar surface.   To guarantee that the disk material close to a luminous and hot supergiant star is neutral, the disk has to be massive (e.g., Kraus \\& Lamers \\cite{KL}). Because the disks around supergiants cannot be pre-main sequence in origin, their formation must be linked to non-spherical (predominantely equatorial), high-density mass loss from the central star, either by some continuous steady material outflow, or by some mass-ejection event(s).   The formation of this anisotropic, high-density mass loss is assumed to be linked to rapid stellar rotation (e.g. Maeder \\& Meynet \\cite{MaMe00}; Kraus \\cite{Kraus06}). Especially rotation close to the critical velocity could trigger the (rotationally induced) bi-stability mechanism (Lamers \\& Pauldrach \\cite{LP}; Pelupessy et al. \\cite{Pelupessy}; Cur\\'{e} \\cite{Cure04}; Cur\\'{e} et al. \\cite{Cure05}) as a plausible disk-formation scenario. And indeed, for at least four members of the B[e] supergiant group, the projected rotation speed (i.e., $\\varv \\sin i$) could be determined, showing that two of them are rotating at a substantial fraction (i.e., at least 75\\%) of their critical velocity. These are the SMC stars \\object{LHA 115-S 23} discussed in detail by Kraus et al. (\\cite{Kraus08}) and \\object{LHA 115-S 65} (Zickgraf \\cite{Zickgraf00}). Whether all B[e] supergiants are rapidly rotating, is unknown however. The high-density wind observed around the other B[e] supergiants hides the central star, and the circumstellar material additionally contributes and pollutes the optical spectrum with a huge amount of emission lines. Consequently, the detection of uncontaminated photospheric lines, which would be appropriate for the determination of the projected stellar rotation velocity, is strongly hampered.  \\begin{table} \\begin{center} \\caption{Parameters of S\\,65. The rotation velocity was derived by Zickgraf (\\cite{Zickgraf00}), the remaining parameters by Zickgraf et al. (\\cite{Zickgraf86}).} \\label{tab1} \\begin{tabular}{cccccc}\\hline \\hline $T_{\\rm eff}$ & $R_{*}$ & $\\log L_*/L_{\\odot}$ & $\\varv\\sin i$ & $E(B-V)$ & inclination \\\\ $ $ [K] &  [$R_{\\odot}$] &   & [km\\,s$^{-1}$] & &  \\\\ \\hline 17\\,000 & 81 & 5.7 & $\\sim 150$ & $0.18\\pm 0.02$ & $\\pm$ edge-on \\\\ \\hline \\end{tabular} \\end{center} \\end{table}   We study here the neutral material around the rapidly rotating SMC B[e] star LHA\\,115-S\\,65 (in short: S\\,65). The star S\\,65 is one of the four confirmed B[e] supergiants in the SMC. Its stellar parameters are summarized in Table\\,\\ref{tab1} and have been determined by Zickgraf et al. (\\cite{Zickgraf86}) and Zickgraf (\\cite{Zickgraf00}). Our investigation makes use of the three optical [O{\\sc i}] emission lines detected in our spectra. We particularly aim to simultaneously reproduce the line luminosities and line profiles, to test the validity of any of the two most plausible B[e] supergiant disk scenarios: (i) an outflowing disk, which is neutral in hydrogen straight from the stellar surface, and (ii) a Keplerian rotating disk. ",
        "conclusions": "We studied the circumstellar material of the rapidly rotating SMC B[e] supergiant star S\\,65. The sharp absorption components seen in its numerous emission lines indicate a high-density disk viewed almost edge-on. Surprisingly, the spectrum also displays two very intense and clearly double-peaked [Ca{\\sc ii}] lines, supporting the high-density gaseous disk hypothesis. Based on the modeling of both, the observed line luminosities of the [O{\\sc i}] lines and their line profiles, we found that the disk around the B[e] supergiant S\\,65 must be predominantly neutral in hydrogen. The total amount of ionized hydrogen is found to be less than $\\sim 0.1$\\%, and the total amount of gas traced within the [O{\\sc i}] line-forming region is $\\sim 1.5\\times 10^{-2}\\,M_{\\odot}$. In combination with the kinematical information provided by the numerous lines, our model results favor the interpretation of a detached (quasi-)Keplerian rotating disk rather than an outflowing disk. The location of S\\,65 in the HR diagram places the star close to the recently defined LBV minimum instability strip. Together with its high rotation velocity close to the critical limit and its large amount of slowly expanding circumstellar material similar to that seen for LBV nebulae, it seems that S\\,65 might be the link between B[e] supergiants and LBVs."
    },
    "1005/1005.3331_arXiv.txt": {
        "abstract": "This paper analyzes a sample of 489 {\\it Spitzer\\/}/IRAC sources in the Extended Groth Strip whose spectral energy distributions fit a red power law from 3.6 to 8.0~\\micron.  The median for sources with known redshift is $\\langle z\\rangle =1.6$. Though all or nearly all of the sample are likely to be active galactic nuclei, only 33\\% were detected in the EGS X-ray survey (AEGIS-X) using 200~ks {\\it Chandra\\/}\\ observations.  The detected sources are X-ray luminous with $L_X > 10^{43}$ erg~s$^{-1}$ and moderately to heavily obscured with $N_H > 10^{22}$~cm$^{-2}$. Stacking the X-ray-undetected sample members yields a statistically significant X-ray signal, suggesting that they are on average more distant or more obscured than sources with X-ray detections.  The ratio of X-ray to mid-infrared fluxes suggests that a substantial fraction of the sources undetected in X-rays are obscured at the Compton-thick level, in contrast to the X-ray-detected sources, all of which appear to be Compton-thin. For the X-ray-detected power-law sources with redshifts, an X-ray luminosity $L_X \\sim 10^{44}$~erg~s$^{-1}$ marks a transition between low-luminosity, blue sources dominated by the host galaxy to high-luminosity, red power-law sources dominated by nuclear activity.  X-ray-to-optical ratios, infrared variability, and 24~\\micron\\ properties of the sample are consistent with the identification of infrared power-law sources as active nuclei, but a rough estimate is that only 22\\% of AGNs are selected by the power law criteria. Comparison of the power-law selection technique and various IRAC color criteria for identifying AGNs confirms that high-redshift samples selected via simple IRAC colors may be heavily contaminated by starlight-dominated objects. ",
        "introduction": "A  census of active galactic nuclei (AGNs) is a necessary prerequisite for constraining the accretion history of the Universe, describing the origin of the cosmic background, and understanding galaxy formation. X-ray detection is often used as an efficient method of finding relatively uncontaminated samples of AGNs \\citep{mus04}.  Large numbers of AGNs can be found in deep X-ray surveys, where the AGN density is an order of magnitude higher than that found in deep optical observations via broad emission lines \\citep[e.g.,][]{hal00, ste02, bau04}. X-ray surveys have now resolved $\\sim$90\\% of the cosmic X-ray background (CXRB) at 2--8~keV \\citep[e.g.,][]{mus00, bau04, hic06}. However, current X-ray telescopes lose sensitivity toward higher energies.  Only 60\\% of the X-ray background is resolved at 6--8~keV, and this fraction drops further to 50\\% at $>$8~keV \\citep{wor04, wor05}.  Synthesis models of the hard X-ray background spectrum require larger numbers of AGNs than are currently observed. Heavily obscured, Compton-thick AGNs are expected to outnumber unobscured AGNs by a ratio of $\\sim$4:1 \\citep[e.g.,][]{mai95, com01, gil01, gil07}, though the exact value is still not well constrained. Deep observations can sometimes detect the most luminous of these hard X-ray sources \\citep{toz06, geo07}, but even the deepest X-ray surveys can miss large populations of obscured AGNs.   With the launch of highly sensitive infrared (IR) instruments aboard the {\\it Spitzer Space Telescope\\/}\\ \\citep{wer04}, IR surveys have become a useful complement to X-ray and visible-light AGN searches.  AGNs are thought to have a dusty environment (often imagined as a torus) surrounding an optically- and X-ray-bright accretion disk centered on a massive black hole (e.g., \\citealt{Antonucci1993,Urry1995}).  Incident X-rays from the central engine can be absorbed and reprocessed by the circumnuclear dust and re-emitted at longer wavelengths, making AGNs IR-bright \\citep{bar87, gra94, nen02}.  The infrared bands are also much less affected by dust extinction than the optical, making them ideal for identifying obscured sources.   There have been several studies of AGN selection using {\\it Spitzer\\/}\\ mid-IR colors \\citep[e.g.,][]{lac04, ste05, hat05} or other IR properties \\citep[e.g.,][]{mar05, don05, dad07, fio08,par08}.  One straightforward way to search for obscured AGNs is to study their spectral energy distributions (SEDs) \\citep{alo06, bar06, don07, car08, don08}.  In the mid-IR, AGNs can often be distinguished by their red power-law SEDs, which can extend from the UV all the way into the far-IR \\citep{ree69, neu79,war87, elv94,Assef2010}.  This behavior is thought to arise from a combination of thermal and non-thermal components near the dusty nucleus that together produce a broad  SED rather than from a single power-law emission source \\citep{rie81,bar87,Kotilainen1992,Klaas2001}. Analysis of broadband SEDs therefore proves to be a powerful diagnostic, differentiating AGNs from sources dominated by starlight. The latter generally have a blue mid-IR power-law SED corresponding to the Rayleigh-Jeans tail of the blackbody spectrum or else do not fit a power law at all due to their 8~\\micron\\ PAH emission or 1.6~\\micron\\ stellar bump.   Finding large samples of AGNs in multiwavelength surveys requires both wide area and good sensitivity.  This paper uses deep {\\it Spitzer Space Telescope\\/}\\ data to select a sample of 489 red power-law sources in the Extended Groth Strip (EGS).  The EGS ($\\alpha = 14^h17^m$, $\\delta = +52\\arcdeg30\\arcmin$) is a premier extragalactic survey field observable from the Northern Hemisphere that fits both criteria \\citep{dav07}.  The EGS has been surveyed at nearly every available wavelength to considerable depth and area, and it spans more than $\\sim$2\\arcdeg\\ by 10\\arcmin\\ in the 9.1~ks IRAC pointings \\citep{bar08}. The region's high ecliptic and Galactic latitudes make it well-suited for studying extragalactic sources with little Galactic contamination.   The unique combination of depth and size makes the EGS an ideal field for an AGN search: it is  shallower but larger than the combined GOODS fields and smaller but deeper than the COSMOS field at most wavelengths. Previous IRAC power-law studies have been done in fields with much deeper X-ray exposures ($\\sim$1~Ms) but shallower IRAC observations \\citep[$\\sim$500~s, e.g.,][]{don07}.  The size and depth of the EGS field in the IR bands allows us to probe different redshift regimes with larger samples and improved statistics.  This study therefore provides a view of power-law sources that is complementary to previous work and extends the work of \\citet{bar06}, who analyzed the IR properties of X-ray selected AGNs in a small portion of the EGS.  Multiwavelength data are utilized to further study the properties of these AGN candidates, including deep \\emph{Chandra X-ray Observatory} pointings to confirm AGN status and control contamination.    This paper proceeds as follows: \\S2 introduces the data sets used, and  \\S 3 describes the power-law AGN selection process.  \\S4 compares the power-law selection with other IR selection methods. \\S5 discusses the sample redshift distribution and \\S6 the  spectral energy distributions.  \\S7 details the X-ray, IR, visible, and radio properties of the power law sample.    A summary is given in \\S8. Throughout this paper, we adopt $\\Lambda$CDM cosmology with $H_0 = 71$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_M = 0.27$, and $\\Omega_{\\Lambda} = 0.73$.  All magnitudes are in AB units unless otherwise noted. ",
        "conclusions": "The IRAC power-law selection technique is designed to identify sources that are dominated by AGN emission in the infrared bands regardless of obscuration.  Only 4\\% of the IRAC sources and 2\\% of the IRAC/MIPS sources in the EGS are identified as power-law AGNs using our criteria.  PL AGNs have SEDs that generally follow AGN templates with either a rough power-law shape from IRAC to visible wavelengths or a downturn in flux density at the visible wavelengths. This may be due to gas and dust obscuring the visible light from PL AGNs.  PL AGNs by definition exhibit no SED `bumps'.  This is unlike much of the rest of the IRAC parent population, which often show clear stellar emission features in their SEDs.  PL AGNs typically have low visible flux densities but are highly luminous overall.   Most of the luminous X-ray sources ($L_{\\rm x} \\ge 10^{43}$ erg~s$^{-1}$) in the EGS are PL AGNs. The high luminosity of PL AGNs allows their detection  at high redshifts, and indeed the EGS PL AGNs lie at a median redshift of $z \\sim 1.6$, much higher than the median redshift of $z \\sim 0.7$ of the parent IRAC population.  One-third of the PL AGN sample have X-ray counterparts, a detection rate greater than the 6\\% found for the parent population.  The detection fraction rises slightly if low-significance X-ray detections are included.  However, among all X-ray-detected IRAC sources, only about one-fifth exhibit red power-law SEDs, while a similar fraction exhibit blue power-law SEDs (with spectral index $\\alpha_{i} > -0.5$); the rest are not well fit by a power law.  As most of the X-ray sources in our field are expected to be AGNs, the lack of a red power law in the IRAC bands is not an indication of inactivity in a galaxy.  Rather, the power-law method finds energetically dominant AGNs and leaves out sources with host galaxy stellar contributions rivaling or exceeding the AGN infrared light. There is a clear division between red power-law sources and blue power-law sources near the break luminosity in the X-ray luminosity function at $L_{X} \\sim 10^{44}$ erg~s$^{-1}$, suggesting a physical transition from sources dominated by host galaxy light to those dominated by nuclear emission.  While fewer than half the PL AGNs are invididually detectable in X-rays with 200~ks {\\it Chandra\\/}\\ exposures, an X-ray stacking analysis found a significant X-ray signal for the rest.  Among the X-ray-detected sources with secure redshifts, most of the sources are obscured with $N_H> 10^{22}$~cm$^{-2}$, but none is likely to be Compton-thick with $N_H > 10^{24}$~cm$^{-2}$.  Assuming the stacked source has a redshift equal to the median, the observed hardness ratio corresponds to a column density of $10^{23}$~cm$^{-2}$, but the ratio of mid-infrared to X-ray fluxes for the stacked source implies that a substantial fraction of the undetected sources are Compton thick.  Eighty-two percent of the PL AGN sample have 24~\\micron\\ counterparts.  Though the AGN fraction in other samples increases with higher 24~\\micron\\ flux densities, the majority of our PL AGNs are not especially mid-IR luminous with 57\\% of them having $S_{24} <100~\\mu$Jy. The  IRAC power law cannot be extended to predict observed 24~\\micron\\ flux density --- only 11\\% of PL AGNs have IRAC power laws that extrapolate to 24~\\micron\\ flux densities  within 10\\% of the actual values.  In two epochs separated by six months, there are only mild indications of infrared flux variability in the PL AGN sample. While about one-third of the PL AGNs varied by at least $4\\sigma$ in at least one IRAC band, only a few objects exhibited variability exceeding $4\\sigma$ across all four IRAC bandpasses. These results are consistent with other studies of AGN variability, but the number of epochs is too small to make definitive conclusions about infrared variability in particular.  Though the power-law selection tends to identify optically-faint galaxies, 82\\% of the PL AGNs have $R$-band detections.  All sources with both $R$-band and hard X-ray detections  have X-ray/optical ratios of 0.1 or greater, consistent with values expected for AGNs.  Nine percent of the PL AGNs have 1.4~GHz counterparts.  Half of these have $S_{1.4 \\rm~GHz} < 100~\\mu$Jy, suggesting that PL AGNs are not preferentially radio-bright. Nearly forty percent of the PL AGNs with radio counterparts have $q\\equiv\\log(S_{24} /S_{1.4\\rm~GHz}) < 0$. These sources have radio-to-mid-infrared flux ratios higher than typical for star-forming galaxies and can be characterized as `radio-loud'; however they make up only a small fraction of the overall PL AGN sample.   Ultimately, to uncover the mysteries of galaxy formation and evolution, a complete and reliable census of AGNs must be obtained for study.  To develop such a census, we must discover what types of sources are identified by all of the various AGN selection techniques.  The power-law selection method returns highly luminous, often obscured, nuclear-dominated galaxies at intermediate to high redshifts.  The method is not intended to select all AGNs or even all infrared AGNs; a rough estimate of the selection fraction is 22\\%. Sources with dominant contributions from stellar emission, though they may comprise a large fraction of the total AGN population, will be excluded from this survey.  However, the power-law technique is intended to reliably select active sources with little contamination, and in particular can select heavily obscured AGNs (and hence AGNs missed by X-ray, visible, or radio surveys). Thus, it provides an excellent way of supplementing other AGN selection methods and leads to a more complete AGN census for study."
    },
    "1005/1005.4727_arXiv.txt": {
        "abstract": "We present the results of near-infrared multi-object spectroscopic observations for 37 $BzK$-color-selected star-forming galaxies conducted with MOIRCS on the Subaru Telescope. The sample is drawn from the $K_s$-band selected catalog of the MOIRCS Deep Survey (MODS) in the GOODS-N region. About half of our samples are selected from the publicly available 24\\,$\\mu$m-source catalog of the Multiband Imaging Photometer for {\\it Spitzer} on board the {\\it Spitzer Space Telescope}. H$\\alpha$ emission lines are detected from 23 galaxies, of which the median redshift is 2.12. We derived the star formation rates (SFRs) from extinction-corrected H$\\alpha$ luminosities. The extinction correction is estimated from the SED fitting of multi-band photometric data covering UV to near-infrared wavelengths. The Balmer decrement of the stacked emission lines shows that the amount of extinction for the ionized gas is larger than that for the stellar continuum. From a comparison of the extinction corrected H$\\alpha$ luminosity and other SFR indicators we found that the relation between the dust properties of stellar continuum and ionized gas is different depending on the intrinsic SFR (differential extinction). We compared SFRs estimated from extinction corrected H$\\alpha$ luminosities with stellar masses estimated from SED fitting. The comparison shows no correlation between SFR and stellar mass. Some galaxies with stellar mass smaller than $\\sim10^{10}\\,{\\rm M}_\\odot$ show SFRs higher than $\\sim100\\,{\\rm M}_\\odot\\,{\\rm yr}^{-1}$. The specific SFRs (SSFRs) of these galaxies are remarkably high; galaxies which have SSFR higher than $\\sim10^{-8}\\,{\\rm yr}^{-1}$ are found in 8 of the present sample. From the best-fit parameters of SED fitting for these high SSFR galaxies, we find that the average age of the stellar population is younger than 100\\,Myr, which is consistent with the implied high SSFR. The large SFR implies the possibility that the high SSFR galaxies significantly contribute to the cosmic SFR density of the universe at $z\\sim2$. When we apply the larger extinction correction for the ionized gas or the differential extinction correction, the total SFR density estimated from the H$\\alpha$ emission line galaxies is $0.089-0.136\\,{\\rm M}_\\odot{\\rm yr}^{-1}{\\rm Mpc}^{-3}$, which is consistent with the total SFR densities in the literature. The metallicity of the high-SSFR galaxies, which is estimated from the N2 index, is larger than that expected from the mass-metallicity relation of UV-selected galaxies at $z\\sim2$ by \\citet{Erb2006a}. ",
        "introduction": "\\label{sec:intro}  Recent studies revealed that a significant fraction of stars in the present-day galaxies were formed between redshifts 1 and 3 \\citep[e.g.,][]{Hopkins2004,PerezGonzalez2005,Hopkins2006,Wang2006,Dahlen2007,Caputi2007}. For example, results of deep near-infrared observations show the average stellar mass density at $0.5<z<1$ is $\\sim 53-72\\%$ of that in the present universe, while at $2.5<z<3.5$ it is just $\\sim4-9\\%$ of the present value \\citep{Kajisawa2009}. Although such studies demonstrate the total star formation rate (SFR) density in the universe as a function of cosmic time, it is still uncertain which populations of galaxies contribute to the active star formation between redshift 1 and 3. The evolution of the stellar mass-SFR relation constrains theoretical views of how galaxies accumulate the stellar mass \\citep{Dave2008}. Therefore, as a next step, in order to understand the build up of stellar mass in galaxies at this redshift range, it is crucial to examine the relationship of SFR to stellar mass.  SFRs are often estimated from photometric indicators that are closely related to the emission from massive stars \\citep{Kennicutt1998}. Massive stars have short lifetimes, so they represent stars which were born recently. The UV continuum (1500--2800\\,\\AA) directly reflects the number of massive stars, and therefore can be used as an indicator of SFR for galaxies in this redshift range. Additionally, far-infrared emission is another indicator of SFR; the UV emission of young massive stars is absorbed by interstellar dust and reradiated in the far-infrared at $\\sim$8--1000\\,$\\mu$m. Fluxes at 24\\,$\\mu$m wavelength taken with the Multiband Imaging Photometer for {\\it Spitzer} (MIPS) are commonly used to estimate far-infrared emission of galaxies at high redshifts, because observations in longer wavelengths are too shallow to detect such galaxies.  \\citet{Daddi2007a} showed that $K$-selected galaxies at $z\\sim2$ have a tight correlation ($\\sim0.2$\\,dex) between stellar mass and SFR estimated from UV and IR luminosities. The SFRs of these galaxies are roughly proportional to their stellar masses. However, the relation is not still conclusive. Other studies show different relationships between SFRs and stellar masses. For example, rest-UV selected BM/BX galaxies also have a correlation between these values but the correlation is not proportional \\citep{Reddy2006a}. Furthermore, the SFRs of 24\\,$\\mu$m-selected galaxies do not depend on their stellar masses \\citep{Caputi2006a}. We need to note that SFRs estimated from UV and IR luminosities have large uncertainties. UV light is affected by dust extinction and stars with long lifetimes of $\\sim1$\\,Gyr. On the other hand, 24\\,$\\mu$m radiation corresponds to the rest-frame 8\\,$\\mu$m wavelength for galaxies at $z\\sim2$. It is known that the far-infrared luminosity estimated from rest-8\\,$\\mu$m luminosity has a large uncertainty due to the complex SED at those wavelengths \\citep{Peeters2004,Chary2001,Daddi2007a,Papovich2007}.  H$\\alpha$ luminosity is one of the most reliable indicators of SFR. First, the emission line is reemitted from hydrogen gas ionized by UV photons shorter than 912\\,\\AA, which are radiated by only the most massive stars. These stars have shorter lifetimes than stars emitting the UV continuum above 1500\\,\\AA. Second, the emission line is less affected by dust extinction than UV emission. The amount of dust extinction is larger at shorter wavelengths. Third, the H$\\alpha$ luminosity can be measured without the uncertainty inherent in SED estimates for the far-infrared luminosity. However, it is difficult to measure H$\\alpha$ luminosities of galaxies at high redshift, because the H$\\alpha$ emission line at $z\\gtrsim0.4$ is out of the wavelength range of optical spectrographs.  Recently, near-infrared spectroscopic studies of SFRs based on the H$\\alpha$ luminosity for large samples of star-forming galaxies at $z\\sim2$ have been conducted \\citep{Erb2006c,Erb2003,Hayashi2009,ForsterSchreiber2009}. \\citet{Erb2006c} demonstrate a weak correlation between SFR estimated from H$\\alpha$ luminosity and stellar mass. Their sample galaxies are selected by $U_nG\\cal{R}$ color and have redshifts confirmed using a blue-sensitive optical spectrograph. The observed galaxies are biased toward less attenuated galaxies. Since red galaxies account for the significant fraction of stellar mass density at $z\\sim2$ \\citep{PerezGonzalez2008}, the galaxies at $z\\sim2$ should be re-examined with a technique more sensitive to red galaxies (e.g., $K$-selected galaxies).  The $K$-selected star-forming galaxies observed by \\citet{Hayashi2009} showed a flatter SFR-stellar mass correlation than that of \\citet{Daddi2007a}, but the number of $K$-selected galaxies with H$\\alpha$ detection is still rather limited and observations for a larger sample are necessary.  In this work, we study the SFR and stellar mass of a large number of $K$-selected star-forming galaxies at $z\\sim2$ based on near-infrared observations with the Multi-Object InfraRed Camera and Spectrograph \\citep[MOIRCS;][]{Suzuki2008}. MOIRCS has the capability for multi-object near-infrared spectroscopy. We also utilize the deep near-infrared imaging of the MOIRCS Deep Survey \\citep[MODS;][]{Kajisawa2009}, which covers $\\sim 103.3\\,{\\rm arcmin}^2$ in the Great Observatories Origins Deep Survey North (GOODS-N) region. The sample for the present work is selected from $K$-selected galaxies in the MODS catalog and their properties are studied using MODS and publicly available multi-wavelength data of GOODS-N. Throughout the paper, we use the AB magnitude system and adopt cosmological parameters of $h\\equiv H_0/100 = 0.7$, $\\Omega_M = 0.3$, $\\Omega_\\Lambda = 0.7$. ",
        "conclusions": "\\label{sec:sum} We performed near-infrared spectroscopic observations for 37 $BzK$ color selected star-forming galaxies with MOIRCS on the Subaru Telescope. The sample galaxies are picked from the deep $K_s$-selected catalog of MODS. We select 18 of our sample from the publicly available 24\\,$\\mu$m source catalog of MIPS on board the $Spitzer$ Space Telescope. In order to avoid contamination by AGNs, we exclude the BzKs detected in the {\\it Chandra} Deep Field-North 2\\,Ms Point-Source Catalogs. All of the sample galaxies have strong 1.6\\,$\\mu$m bumps in their SEDs, so that their optical emission is thought to be dominated by the stellar continuum. The H$\\alpha$ emission lines are detected from 23 galaxies, of which the median redshift is $z=2.12$. In the diagnostic diagram with \\othree$\\lambda$5007/H$\\beta$ and \\ntwo$\\lambda$6583/H$\\alpha$, the line ratios of the stacked spectra are above that found for local star-forming galaxies but closer to the expectations from star-formation activity rather than AGN activity.  We derived the SFRs from the extinction-corrected H$\\alpha$ luminosities. The extinction correction for the H$\\alpha$ luminosity is estimated from the SED fitting of multi-band photometric data with stellar population models of exponentially declining SFR and the \\citet{Calzetti2000} extinction curve. To estimate the extinction correction for the ionized gas, we assumed the \\citet{Calzetti2001} relation ($E(B-V)_{\\rm stellar} = 0.44\\times E(B-V)_{\\rm gas}$) as well as the equal-extinction case $E(B-V)_{\\rm stellar} = E(B-V)_{\\rm gas}$, and compared the extinction corrected SFRs. Although the Balmer decrement of the stacked emission lines is consistent with the \\citet{Calzetti2001} relation, a fraction of the H$\\beta$ detected galaxies are consistent with the equal-extinction case. We further investigated these galaxies by comparing the extinction corrected H$\\alpha$ luminosities with other SFR indicators. The comparison with the UV and MIPS 24\\,$\\mu$m flux shows that the SFRs of lower SFR galaxies ($\\lesssim100\\,{\\rm M}_\\odot\\,{\\rm yr}^{-1}$) are in good agreement for the equal-extinction case, while those of the higher SFR galaxies agree for the \\citet{Calzetti2001} relation. This implies that the relation between the dust properties of the stellar continuum and emission lines are different depending on the intrinsic SFR (differential extinction correction). If the Calzetti relation is assumed, a fraction of the high-SFR galaxies should be detected in radio, (sub-)mm, and X-ray catalogs, but none of them except one AzTEC source are detected. We could not explain the reason for this discrepancy consistently. However, most of the discrepancy could be explained by the short timescale of star-formation activity, because the younger ages of the stellar population are derived from our SED analysis. We note that other mechanisms (e.g., top-heavy IMF, lower dust temperature, etc.) could also explain a part of the discrepancy.  We compared the extinction-corrected H$\\alpha$ luminosities with the stellar masses obtained from the SED fitting. When we apply the \\citet{Calzetti2001} relation, the comparison shows no correlation between SFR and stellar mass. A fraction of the galaxies with stellar mass smaller than $\\sim10^{10}\\,{\\rm M}_\\odot$ show SFRs higher than $\\sim100\\,{\\rm M}_\\odot\\,{\\rm yr}^{-1}$. The SSFRs of these galaxies are higher than $\\sim10^{-8}\\,{\\rm yr}^{-1}$. Even if equal extinction is applied, their SSFRs are still high (${\\rm SSFR}>3.5\\times10^{-9}\\,{\\rm yr}^{-1}$). This is inconsistent with the correlation demonstrated by \\citet{Daddi2007a}. Our value of extinction correction shows large scatter against $B-z$ color and stellar mass, with which \\citet{Daddi2004b} and \\citet{Pannella2009} show strong correlations. From the best-fit parameters of the SED fitting for these high-SSFR galaxies, we found that the average age of the stellar population is younger than 100\\,Myr, which are consistent with SFR age (${\\rm SSFR}^{-1}$) with the \\citet{Calzetti2001} relation. We also calculate the contribution of the present H$\\alpha$ emission to the cosmic SFR density at $z\\sim2$. Our results are consistent with previous works, if we apply the \\citet{Calzetti2001} relation or differential extinction correction. The high-SSFR galaxies significantly contribute ($50-75\\%$) to the cosmic SFR density.  The metallicity of the high-SSFR galaxies, which is estimated from the N2 index, is larger than that expected from the mass-metallicity relation of UV-selected galaxies at $z\\sim2$ by \\citet{Erb2006a}. Although metallicities estimated from the N2 index is known to have significant uncertainties due to the harder ionizing radiation field and ionization parameter, the high metallicity comparable to galaxies in the same mass range is inconsistent with the known star-formation history of such galaxies in the local universe, which formed a significant fraction of their stars in the low redshift universe. The high-SSFR galaxies, therefore, possibly evolve into larger stellar mass galaxies with low star-forming efficiency, depleting the gas."
    },
    "1005/1005.1116_arXiv.txt": {
        "abstract": "The invariant four-volume $\\mathcal{V}$ of a complete black hole (the volume of the spacetime at and interior to the horizon) diverges. However, if one considers the black hole set up by the gravitational collapse of an object, and integrates only a finite time to the future of the collapse, the resultant volume is well defined and finite. In this paper we examine non-degenerate stationary black holes (and cosmological horizons) and find that  $\\mathcal{V}_{s} \\varpropto \\ln(\\lambda)$ where $s$ is any  shell that terminates on the horizon, $\\lambda$ is the affine generator of the horizon and the constant of proportionality is the Parikh volume of $s$ divided by the surface gravity. This provides an alternative local and invariant definition of the surface gravity of a stationary black hole. ",
        "introduction": "The purpose of this paper is to discuss the importance of the rate of growth of the invariant four-volume $\\mathcal{V}$ of stationary non-degenerate black holes and how this rate of growth is related to their surface gravity. Now whereas the three-volume of a black hole depends on the choice of slicing \\cite{din}, and in general the associated three-volumes are finite, the full four-volume $\\mathcal{V}$ is usually never discussed since it is formally infinite. (Consider, for example, the four-volume of the region on and below the future horizons in the Kruskal-Szekeres plane.) Even if we consider the black hole as created by the gravitational collapse of an object and so consider only part of the Kruskal-Szekeres plane on and below the future horizon and to the future of the boundary surface of the collapsing object, the four-volume still diverges as we integrate to the infinite future. However, we need not integrate to the infinite future and so in effect consider the evolution of $\\mathcal{V}$. In this paper, to start, we examine explicitly in regular coordinates the four-volume bounded by the horizon, the central singularity and two distinct ingoing null cones in the Schwarzschild spacetime. The situation is shown schematically in Figure \\ref{figure}. This introductory calculation points out a relation which we eventually show is a universal feature of stationary black holes: The invariant 4-volume $\\mathcal{V}$ of a black hole grows as  $\\mathcal{V} \\varpropto \\ln(\\lambda)$ where $\\lambda$ is the affine generator of the horizon and the constant of proportionality is the Parikh volume $^3\\mathcal{V}$ \\cite{parikh} (equivalent to the Euclidean 3-volume in spherical symmetry) divided by the surface gravity ($\\kappa$). Defining $ \\mathcal{V}^{*} \\equiv d \\mathcal{V} /d \\ln(\\lambda)$ then we have an alternative definition of the surface gravity: $\\kappa=\\;^3\\mathcal{V}/\\mathcal{V}^{*}$. However, the Reissner-Nordstr\\\"{o}m-de Sitter class of black holes, for example, contain unstable Cauchy horizons within their event horizons \\cite{cauchy} and so one might argue that in such cases we would be integrating over unphysical regions of the spacetime. To overcome this we show that the four-volume of a shell ($s$) which terminates on the horizon is well defined and the shell can obviously be constructed so as to exclude the Cauchy horizons. In this case we find $\\kappa=\\;^3\\mathcal{V}_{s}/\\mathcal{V}^{*}_{s}$ and so $\\kappa$ is in fact independent of the ``depth\" of integration. We use such a construction to consider cosmological horizons and to argue that because of this independence on depth, $\\kappa$ is a property of the horizon alone.  \\bigskip  The paper is organized as follows: First we give a general overview of the static spacetimes involved with our preliminary considerations. Then the invariant partial 4-volume of a Schwarzschild black hole is calculated explicitly in regular coordinates. To consider more general cases we first review the construction of regular coordinates for non-degenerate static black holes. We then calculate the invariant 4-volume $\\mathcal{V}$ in this more general context and explain the shell construction necessary for the consideration of cosmological horizons. To move on to the stationary case we first review Parikh's construction and we then show that $\\kappa=\\;^3\\mathcal{V}_{s}/\\mathcal{V}^{*}_{s}$ holds in general. The explicit verification in the Kerr metric is relegated to an Appendix. We conclude with a discussion of this alternative definition of the surface gravity. \\begin{figure}[ht] \\epsfig{file=blackholer.eps,height=3in,width=2.7in,angle=0} \\caption{\\label{figure}The collapse of a timelike boundary surface $\\Sigma$ that terminates at the central singularity simultaneously with the null cone $\\delta$ and produces a black hole with horizon $\\mathcal{H}$. The null cone $v$ is any null cone to the future of $\\delta$. The partial invariant four-volume $\\mathcal{V}$ calculated here is bounded by $\\delta$ and $v$, and is on and to the interior of $\\mathcal{H}$.  We eventually show that the depth of the integration below $\\mathcal{H}$ is irrelevant as regards the surface gravity $\\kappa$. } \\end{figure} ",
        "conclusions": "The usual physical meaning given to the surface gravity is, as explained for example by Poisson \\cite{cauchy}, ``the force required of an observer at infinity to hold a particle (of unit mass) stationary at the horizon\" (think of the Schwarzschild case). The interpretation given here, that the surface gravity is the ratio of the Parikh volume to the rate of change of the invariant four-volume for a shell of arbitrary (but non-vanishing) thickness bounded by the horizon, is a local interpretation that would appear to be new. An important question is, can we use this to gain further insights into black hole mechanics? The most obvious conclusion regards the third law of black hole mechanics, $\\kappa \\nrightarrow 0$. Since $^3\\mathcal{V}_{s}>0$ even in the degenerate case (consider, for example, the static spherically symmetric case) we see that the third law demands that the rate of growth $\\mathcal{V}^{*}_{s}$ must remain finite. In order to violate the third law we need $\\mathcal{V}^{*}_{s} \\rightarrow \\infty$ and since $d\\mathcal{V}_{s}/d \\lambda$ is finite, this requires $\\lambda \\rightarrow \\infty$ in agreement with the formulation of Israel \\cite{israel}. That is, in a sequence of quasi-static steps, the reduction of $\\kappa$ to zero would take infinite advanced time.  \\bigskip  \\textit{Acknowledgments.} It is a pleasure to thank Maulik Parikh who pointed out \\cite{parikh} the content of which changed our presentation from an earlier draft. This work was supported by a grant to KL from the Natural Sciences and Engineering Research Council of Canada. Portions of this work were made possible by use of \\textit{GRTensorII} \\cite{grt}.  \\bigskip  \\appendix*"
    },
    "1005/1005.1927_arXiv.txt": {
        "abstract": "A Galileon field is one which obeys a spacetime generalization of the non-relativistic Galilean invariance. Such a field may possess non-canonical kinetic terms, but ghost-free theories with a well-defined Cauchy problem exist, constructed using a finite number of relevant operators. The interactions of this scalar with matter are hidden by the Vainshtein effect, causing the Galileon to become weakly coupled near heavy sources. We revisit estimates of the fifth force mediated by a Galileon field, and show that the parameters of the model are less constrained by experiment that previously supposed. ",
        "introduction": "\\label{sec:intro}  Many disparate cosmological observations suggest that the universe recently entered an accelerating epoch. The simplest explanation for this acceleration is the inclusion of a cosmological constant $\\Lambda \\sim (10^{-3}\\;\\mbox{eV})^4$ in the Einstein equations. Within classical general relativity there is no reason to prefer one value of $\\Lambda$ over another. The inclusion of particle physics, however, leads to a natural interpretation of $\\Lambda$ as the energy density of the vacuum.  Field-theoretic estimates of this energy density diverge as the fourth power of the cutoff, $M$, yielding $\\Lambda \\sim M^4$. A realistic choice almost certainly requires $M \\gg \\mbox{TeV}$, which markedly disagrees with the tiny value measured by experiment.% \\footnote{If the cutoff for field theory is assumed to be the Planck scale, there is a mismatch between the predicted and observed values of $\\Lambda$ by approximately 120 orders of magnitude.}  The ``old'' cosmological constant problem is the question: why is the observed value of the cosmological constant so small \\cite{Weinberg:2000yb}? But we can just as well ask: why does the vacuum energy gravitate so little \\cite{Dvali:2007kt}? If long wavelength fluctuations gravitate more weakly than other sources, the vacuum energy may be screened or ``degravitated'' \\cite{Dvali:2002pe,ArkaniHamed:2002fu} in the same way that long wavelength excitations far beyond the Debye sphere are screened by the effective photon mass in a plasma. By analogy, degravitation occurs if the graviton can acquire a mass, perhaps because it occurs as a resonance \\cite{Dvali:2007kt} or extended object \\cite{Sundrum:2003jq}, or because it descends from a braneworld scenario such as the DGP model or its generalizations \\cite{deRham:2007xp,deRham:2007rw}. An implementation of Higgs' mechanism for the graviton has recently been constructed by Chamseddine \\& Mukhanov \\cite{Chamseddine:2010ub}.  The four-dimensional effective theory of such models necessarily includes extra scalar fields, arising from new longitudinal polarizations of the graviton. It is these new modes which are the agent of degravitation, requiring them to couple to gravitational sources. Therefore one may have some reservations that these scalars mediate unacceptably large fifth forces. However, it is possible for these fields to avoid experimental constraints using the Vainshtein (``kinetic chameleon'') mechanism \\cite{Vainshtein:1972sx,Deffayet:2001uk}, which exploits their non-canonical kinetic terms. Non-linearities become important in the vicinity of dense objects, and the scalars naturally become weakly coupled to matter. A successful Vainshtein decoupling may require less delicate fine-tuning of parameters.  Nicolis, Rattazzi \\& Trincherini \\cite{Nicolis:2008in} discussed the most general four-dimensional ghost-free Lagrangian for a scalar field with these properties. Their construction invoked a gradient shift symmetry for the scalar field, a generalization of non-relativistic Galilei invariance, which permits sizeable non-linearities to occur. Nicolis {\\etal} noticed that a finite set of operators were consistent with this symmetry and simultaneously gave rise to second-order equations of motion. Taking these operators as the relevant terms in an effective field theory, with all other operators considered as irrelevant deformations, one arrives at a perturbatively unitary Lagrangian \\cite{Donoghue:1994dn,Donoghue:1993eb}. Moreover, the relevant operators are protected against renormalization from scalar loops \\cite{Luty:2003vm,Porrati:2004yi}.  Nicolis {\\etal} denoted the field described by the resulting Lagrangian as $\\pi$, and referred to it as a \\emph{Galileon}. Couplings of the Galileon to gravity were studied by Deffayet {\\etal} \\cite{Deffayet:2009wt,Deffayet:2009mn}, and aspects of its cosmological phenomenology were explored by numerous authors \\cite{Chow:2009fm,Silva:2009km,Kobayashi:2010wa,Kobayashi:2009wr,Gannouji:2010au}. The gravitational backreaction of the Galileon field was shown to be small in Ref.~\\cite{VanAcoleyen:2009kb}. An explicit higher dimensional theory, for which the Galileon is the four-dimensional effective theory, was discussed in Ref.~\\cite{deRham:2010eu}.  The Galileon Lagrangian contains at least five dimensionful parameters, one for each of the relevant operators.% \\footnote{Other parameters may be associated with irrelevant deformations, which we ignore in the present discussion.} These may be chosen freely subject to certain constraints. Of these, the most important is that the acceleration of the Hubble rate should agree with our observations. The remaining parameters must combine in such a way that fifth forces mediated by the Galileon are phenomenologically acceptable. These constraints were discussed by Nicolis {\\etal} \\cite{Nicolis:2008in}. In this paper we revisit this analysis and show that the Vainshtein radius of a massive object---within which the equations of motion of the Galileon become non-linear and the scalar force it mediates is dynamically suppressed---% can be made arbitrarily large by varying a single combinations of the parameters.  The magnitude of the scalar force within the Vainshtein radius depends on an independent combination of parameters. When a sufficient number of relevant operators are included, these quantities may therefore vary independently.  This freedom leaves the parameters of the Galileon model weakly constrained by gravitational experiments.  With this in mind we revisit the strong-coupling limit of the theory and show that, in contrast to the DGP scenario, the distance scale at which non-renormalizable corrections become important may be far below the reach of current experimental probes.  This brief note is organized as follows. In \\S\\ref{sec:eft} we review the action for a  Galileon scalar field.  In \\S\\ref{sec:fifth} the fifth forces mediated by Galileon exchange are calculated, and we demonstrate that the constraints on the parameters of the theory are not as stringent as previously supposed.  For this broader parameter space, the scale at which quantum fluctuations become strongly interacting is derived in \\S\\ref{sec:strong}. We conclude in \\S\\ref{sec:conc}. In an Appendix we give a simple derivation of the Galileon force law assumed throughout this article. We choose units in which $c = \\hbar = 1$. ",
        "conclusions": "\\label{sec:conc}  The Galileon is a scalar field coupled to the matter content of the Standard Model. Its effects can be hidden from searches for modifications of gravity, without gross fine-tuning of parameters, because its non-linear self-interactions cause the field to become weakly coupled near heavy objects. The form of these self interactions is protected by a derivative shift symmetry.  Our discussion has been framed in the context of an effective field theory, where the Galileon's behaviour is controlled by a small number of relevant couplings. In some previous discussions, compatibility with a higher-dimensional completion has fixed these couplings and therefore the Vainshtein radius. In Ref.~\\cite{Nicolis:2008in} the couplings were allowed to vary freely, but the far-field interaction strength, controlled by the coupling $d_2$, was set to gravitational strength. In this article we have shown that the Vainshtein radius associated with a massive object can be made arbitrarily large or small by varying a linear combination of the Lagrangian parameters. Experiments probing the force on a test mass interacting with a heavy source at fixed radius can be made insensitive to Galileon exchange forces by varying an independent linear combination of Lagrangian parameters. In principle this allows the far-field interaction strength to be made much stronger. Whether such an arrangement can be technically natural is presently unknown.  Calculation of the precise constraints imposed by experimental tests of gravity is left for future work. However, it seems reasonable that observations do not force us to put all the Lagrangian parameters at the Planck scale, as had been supposed. Instead, they may vary widely over many orders of magnitude.  This is particularly interesting in light of the recently proposed five-dimensional origin of the Galileon model \\cite{deRham:2010eu}, where the Galileon parameters derive from combinations of the parameters of the five-dimensional theory. The freedom to vary the Galileon parameters also means that the quantum strong coupling scale can be made arbitrarily high, and placed out of reach of current experimental probes."
    },
    "1005/1005.4675_arXiv.txt": {
        "abstract": "{The interstellar medium is enriched primarily by matter ejected from evolved low and intermediate mass stars. The outflows from these stars create a circumstellar envelope in which a rich gas-phase and  dust-nucleation chemistry takes place. We observed the nearest carbon-rich evolved star, IRC\\,+10216, using the PACS (55-210\\,$\\mu$m) and SPIRE (194-672\\,$\\mu$m) spectrometers  on board Herschel. We find several tens of lines from SiS and  SiO, including lines from the v=1 vibrational level. For SiS these transitions range up to J=124--123, corresponding to energies around 6700\\,K, while the highest detectable transition is J=90--89 for SiO, which corresponds to an energy around 8400\\,K. Both species trace  the dust formation zone of IRC\\,+10216, and the broad energy ranges involved in their detected transitions permit us to derive the physical properties of the gas and the particular zone in which each species has been formed. This allows us to check the accuracy of chemical thermodynamical equilibrium models and the suggested depletion of SiS and SiO due to accretion onto dust grains.} {} {} {} {} ",
        "introduction": "\\object{IRC\\,+10216} (\\object{CW\\,Leo}) is the brightest non-Solar System object in the sky at 5\\,$\\mu$m.  It is the nearest \\citep[D$\\sim$150\\,pc,][]{Crosas1997ApJ...483..913C} carbon-rich evolved star, and it serves as an archetype for the study of mass loss on the asymptotic giant branch (AGB). The star is losing mass at a rate of $\\sim$1--$3\\times10^{-5}$\\,\\Msun/yr \\citep{Crosas1997ApJ...483..913C, Schoier2000A&A...359..586S}, producing a dense, dusty circumstellar envelope (CSE). To date, more than 60 molecules have been detected in the CSE of IRC\\,+10216 \\citep[e.g.,][]{Cernicharo2000A&AS..142..181C, He2008ApJS..177..275H}. Thermodynamic equilibrium and non-equilibrium reactions, photochemical reactions, ion-molecule reactions, and the condensation of dust grains establish the abundance stratifications throughout the envelope. It is likely that IRC\\,+10216 is  in an advanced evolutionary stage, marking the transition from an AGB to a planetary nebula \\citep{Skinner1998MNRAS.300L..29S}. A detailed study of its infrared emission spectrum can yield unique information on the thermophysical and chemical structure of the outflow and on the history of mass loss during this important evolutionary phase.   \\vspace*{-2ex} ",
        "conclusions": "The PACS and SPIRE spectroscopic observations of  \\irc\\ have been shown to be of excellent quality for studying the thermodynamical and chemical structure of the envelope, created by its copious mass loss. The temperature and mass-loss rate of the envelope are derived from the $^{12}$CO lines. Both SiO and SiS are refractory species, and the PACS and SPIRE data can provide a strong diagnostic tool for determining their role in the dust formation process. Analysing the high-J SiO and SiS lines yields a constant fractional abundance of $1\\times10^{-7}$ and $4\\times10^{-6}$, respectively. However, we detect only v=0 and v=1 transitions for both species, mainly because of the densely populated spectrum of IRC+10216, while it is known from ground-based observations that levels of SiS up to v=8 have been detected (Ag\\'undez et al., 2010, \\emph{in prep.}). Moreover, the low-J transitions of SiO and SiS, which are more sensitive to the external envelope, are not accesible to PACS and SPIRE. Since the high-J lines in the ground-state and the v=1 lines of both molecules are very weak, we cannot put strong constraints on the fractional abundance in the inner envelope (R$\\la10$\\,\\Rstar). For SiO, 1/3 at most is estimated to take part in dust formation process, while we deduce a fraction of 1/2 for SiS. Only a merged set of millimeter, submillimeter, and far-infrared observations of SiO and SiS can provide a detailed analysis of the abundance of these species from the photosphere to the photodissociation zone (Ag\\'undez et al., 2010, \\emph{in prep.}).  \\vspace*{-1ex}"
    },
    "1005/1005.3263_arXiv.txt": {
        "abstract": "During the month of  December, 2009 the blazar 3C~454.3 became the brightest gamma-ray source in the sky, reaching a peak flux $F \\sim 2000 \\times 10^{-8} $ph cm$^{-2}$ s$^{-1}$ for E $>\\ 100$ MeV. Starting in November, 2009 intensive multifrequency campaigns monitored the 3C~454 gamma-ray outburst. Here we report the results of a 2-month campaign involving AGILE, INTEGRAL, \\emph{Swift}/XRT, \\emph{Swift}/BAT, RossiXTE for the high-energy observations, and \\emph{Swift}/UVOT, KANATA, GRT, REM for the near-IR/optical/UV data. The GASP/WEBT provided radio and additional optical data.  We detected a long-term active emission phase lasting $\\sim$1 month at all wavelengths: in the gamma-ray band, peak emission was reached on December 2-3, 2009. Remarkably, this gamma-ray super-flare was not accompanied by  correspondingly intense emission in the optical/UV band that reached a level substantially lower than the previous observations in 2007-2008.  The lack of strong simultaneous optical brightening during the super-flare and the determination of the broad-band spectral evolution severely constrain the theoretical modelling. We find that the pre- and post-flare broad-band  behavior can be explained by a one-zone model involving SSC plus external Compton emission from an accretion disk and a broad-line region. However, the spectra of the Dec. 2-3, 2009 super-flare and of the secondary peak emission on Dec. 9, 2009 cannot be satisfactorily modelled by a simple one-zone model. An additional particle component is most likely active during these states. ",
        "introduction": "The flat spectrum radio quasar 3C 454.3 (at a redshift $z=0.859$) turns out to be among the most active blazars emitting a broad-spectrum ranging from radio to gamma-ray energies. Blazars are a sub-class of active galactic nuclei, with the relativistic jet aligned to the line of sight. Their spectral energy distributions (SED) typically show a double humped shape, with the low energy peak lying between radio and X-rays, and the high energy peak in the GeV-TeV band \\cite{padovani1995}. Detailed description of blazar leptonic emission models can be found in Maraschi et al. (1992); Marsher \\& Bloom (1992); Sikora et al. (1994). The observed spectra can also be modelled in the framework of hadronic models \\cite{mucke2001,mucke2003,bottcher2007}.  Starting in 2004-2005,  3C 454.3 showed a long period of optical activity, with variability timescale ranging from several months to less than one day. In May 2005, the source reached a peak magnitude  $R \\simeq 12$ showing strong 1-day variability \\cite{Villata2006}. A radio peak was detected 9 months after the optical peak, with a flux of 22 Jy at 37 GHz, and  20 Jy at 43 GHz \\cite{villata2007}. The source was then quiescent from the beginning of  2006 until mid-2007 \\citep[with an R magnitude between 15 and 16, ][]{raiteri2008}. Starting in the second half of 2007 \\cite{vercellone2008}, \\c has been detected in a high  gamma-ray state by AGILE \\cite{agile}, and, subsequently also  by  \\emph{Fermi}-LAT \\cite{atwood2009,abdo2009a}. Typically, the level of gamma-ray activity (with a flux of 300-600$\\times 10^{-8}$ ph cm$^{-2}$ s$^{-1}$ for E $>$ 100 MeV) has been observed to be correlated with the optical emission. Relatively large gamma-ray fluxes were detected during the AGILE {\\rffa{observations of 2007 \\cite{vercellone2008,vercellone2009a,donnarumma2009a}.}} During 2008,  the source that was bright at the beginning of the year, started to fade in optical band \\cite{villata2009}. AGILE detected the fading in gamma-rays too \\cite{vercellone2009b}. \\textit{Fermi} reported an averaged spectrum above 200 MeV with a photon index $\\alpha \\simeq 2.3 $ and a spectral break at $E_c \\sim 2.4 $~GeV, obtained in August, 2008 \\cite{abdo2009a}. ",
        "conclusions": "The multifrequency data of the extensive campaign on \\c show a remarkable behavior of the source. Starting from MJD 55150, first an overall rise of the gamma-ray emission  and then of the X-ray and optical fluxes is detected. This rise culminates with peak optical/X-ray/gamma-ray emission during a 10-day period centered around MJD 55173. Subsequently, the overall flux decreased and reached again a relative-minimum state around MJD 55200. During the 2-month period the optical and X-ray fluxes vary within a factor of 3, whereas the \\gray flux grows by a factor of 5-10 compared to the pre-flare value. During the rapid super-flare around MJD~55167.7 the \\gray flux doubles within 1 day with the optical and average X-ray  increase of 50\\% and 30\\%, respectively.  We find an overall correlation at all wavelengths for both long and short timescales. However, the unusual gamma-ray flaring and super-flaring activity from \\c during the period November-December, 2009 is not accompanied by strong emission of similar intensity in the optical or even in the soft X-ray bands. {\\rffa{The hard x-ray flux is comparable to the level of spring 2005, e.g., $\\sim 3.5\\times 10 ^{-10}\\ erg\\ cm^2\\ s^{-1}$ in 20-200 keV \\cite{pian2006}.}} This flaring behavior appears to be quite different from other episodes detected in 2007 and 2008 \\citep[e.g.][]{vercellone2009a,donnarumma2009a}. The synchrotron emission appears to be quite broad and centered around $\\nu \\sim 10^{13}$~Hz.   \\begin{table*} \\begin{center} \\caption{ Models parameters used for producing the spectral analysis of Fig. 2. The columns give the emission interval specification, the leptonic component that dominates the gamma-ray emission, the magnetic field $B$, the comoving emission radius $R$, the particle energy distribution normalization $K$, the break energy (Lorentz factor) $\\gamma_b$, the minimum particle energy $\\gamma_{min}$, and the low- ($\\zeta_1$) and high-energy spectral indices ($\\zeta_2$) defined as $N(\\gamma) = K \\,\\gamma_b^{-1} / [(\\gamma/\\gamma_b)^{\\zeta_1} + (\\gamma/\\gamma_b)^{\\zeta_2}] $. The relativistic Maxwellian is defined as $N(\\gamma) = K \\,(\\gamma/\\gamma_b) \\exp(-\\gamma/\\gamma_b)$.} \\label{table:sed-bi-zone} \\small \\noindent \\scriptsize \\begin{tabular}{|l|c|c|c|c|c|c|c|c|c|c|} \\hline \\bf{Interval}       & \\bf{Model}    &\\bf{Component} &  \\bf{B} & \\bf{R} & $\\bf{K}$    & $\\bf{\\gamma_b}$ & $\\bf{\\gamma_{min}}$ & $\\bf{\\zeta_1}$ & $\\bf{\\zeta_2}$ & Comments \\\\ &               &                          &  (G)    & (cm)   & $(cm^{-3})$ &                 &                     &                &                &          \\\\ \\hline 1 (pre-flare)       & two-comp.     & component-1$^*$   &  0.6 & $7\\times10^{16}$ &2.2   & 800 & 35 & 2.35 &4.5 & broken PL\\\\ \\cline{3-11} &               & component-2   &  --  & --              & --   & --  & -- & --   & -- & --        \\\\ \\hline 2 (super-flare)     & two-comp.     & component-1   &  0.6 & $7\\times10^{16}$ &2.2   & 800 & 35 & 2.35 &4.5 & broken PL\\\\ \\cline{3-11} &               & component-2   &  0.9 & $3\\times10^{16}$ &180   & 180 & 1  &  --  & -- & relativistic Maxwellian  \\\\  \\cline{2-11} &\\multicolumn{2}{|c|}{one-zone} &  0.5 & $7\\times10^{16}$ &2.5   &1000 & 35 & 2.35 & 4.5& broken PL\\\\ \\hline 3 (secondary-flare) & two-comp.     & component-1   &  0.6 & $7\\times10^{16}$ &2.5   & 800 & 45 &2.25 &4.5 & broken PL\\\\ \\cline{3-11} &               & component-2   &  0.9 & $3\\times10^{16}$ &170   & 170 & 1  & --  & -- & relativistic Maxwellian \\\\ \\hline 4 (post-flare)      & two-comp.     & component-1$^*$   &  0.6 & $7\\times10^{16}$ &2.5   & 800 & 45 &2.25 &4.5 & broken PL\\\\ \\cline{3-11} &               & component-2   &  --  & --              & --   & --  & -- & --   & -- & --       \\\\ \\hline \\end{tabular} \\end{center} \\normalsize $(^*)$For the pre- and post-flare intervals this set of parameters describes also the simple one-zone model. \\end{table*}   Striani et al. (2010) provided a first report of the AGILE-GRID data. A single power-law approximation gives a photon spectral index 1.66 $\\pm$ 0.32 in the energy band 0.1-1 GeV, integrating the data for two days centered at MJD 55167.7. We confirm this result in the analysis presented here integrating gamma-ray data for 1 day, as shown in our Fig. 2 (left panel). In addition, we found a similar spectral shape of the gamma-ray emission during the secondary maximum of interval-3. Our Fig. 2 (right panel) shows the interval-3 SED as compared with the post-flare SED of interval-4.  The simultaneous observations of \\c by AGILE, INTEGRAL, RossiXTE and \\textit{Swift} strongly constrain the emission models in the high-energy range. In particular, our interval-3 spectrum (blue solid squares of Fig. 2, right panel) represents one of the best constrained multifrequency spectra ever obtained for a flaring blazar from X-ray up to GeV energies.  We present in Fig. 2 the results of our spectral modelling based on {\\rffa{Synchrotron-Self}} Compton (SSC), {\\rffa{plus contribution from external seed photons (EC).}} We used  parameters similar to those already implemented to model previous gamma-ray flares of 3C~454.3. \\cite{vercellone2009a,donnarumma2009a}. We find that the pre- and post-flare spectra (interval-1 and interval-4) are adequately represented by a simple one-zone SSC model plus EC in which the accretion disk and the broad-line region provide the necessary soft radiation field for the inverse Compton components that dominate the X-ray through the GeV energies. {\\rffa{The simple one-zone model parameters for the pre- and post-flare intervals are reported in Table \\ref{table:sed-bi-zone} (they are also labelled \\emph{component-1} in the table).}}  {\\rffa{We can fit the super-flare with a simple one-zone model as shown in Figure \\ref{fig:sed} (dot-dashed curve), with the parameters reported in Table \\ref{table:sed-bi-zone}. Such a model requires (compared to the pre-flare fit parameters) an increase of the electron energy and density, and a slight reduction of the comoving magnetic field for the whole  electron population of the blob  \\citep[as also reported in ][]{bonnoli2010}.}} {\\rffa{We note, however, that we used the optical data obtained during the rising edge of optical emission for the SED of the super-flare. Using only the V peak emission for the optical portion of the spectra, the one-zone model gives a magnetic field of 0.55 G, closer to the value for the pre-flare.}}  {\\rffa{The time evolution of the source (e.g., the lack of strong optical emission and the gamma-ray spectrum during the super-flare) can be also described adopting a different approach.}} We assume a long-term rise and fall of the mass accretion rate onto the central black hole. This enhanced accretion causes an overall increase of the synchrotron {\\rffa{emission}} and of the soft photon background scattered off by the primary component of accelerated electrons (component-1). An additional population of accelerated leptons (component-2, co-existing with component-1) {\\rffa{can be}} introduced for the super-flare and secondary flare episodes. This component is a consequence of additional particle acceleration and/or plasmoid ejection near the jet basis. Table~1 reports the parameters that we used to model the Fig.~2 spectra {(the two-component models are reported as {\\rffa{thick}} solid lines for the super-flare and for the secondary flare).} {\\rffa{For all intervals, we assume the presence of component-1. This is the only component acting for the pre- and post-flare intervals, with the parameters given in Table~1. The super-flare (secondary-flare) parameters for the component-1 are the same as the pre-flare (post-flare) interval.}} We assumed a bulk Lorentz factor $\\Gamma = 25$, a jet angle with respect with the line of sight $\\theta=1.2^o$, an accretion disk of bolometric luminosity $L_d=6 \\times 10^{46}$\\ergsec slowly decaying toward $L_d= 5 \\times 10^{46}$\\ergsec. A broad line region located 0.5 pc  from the black hole reflects 5\\% of the disk power toward the emitting regions.  The component-2 energy distribution that better reproduces our gamma-ray spectral data is a relativistic Maxwellian of characteristic energy $\\gamma_b \\simeq 180$. Interestingly, this component appears to be strongly energized but not yet modified by additional non-thermal acceleration.  {\\rffa{This two-component approach avoids the problem of explaining the time variability of the physical parameters of the whole electron population of the one-zone model, at the price of adding an additional varying component superimposed with the pre-existing one.}}   To summarize, our multifrequency data for the December, 2009 flare of \\c provide a wealth of very important information on this puzzling and fascinating blazar. We find that \\c is characterized by strong broad-band spectral variability, and that the modelling of the peak gamma-ray emission episodes {\\rffa{suggests more elaborate models than the standard one-zone SSC+EC models of bright blazars.}}  \\bigskip   The AGILE Mission is funded by the Italian Space Agency (ASI) with scientific and programmatic participation by the Italian Institute of Astrophysics (INAF), and the Italian Institute of Nuclear Physics (INFN). This investigation was carried out with partial support from the ASI contract n.~I/089/06/2. V.Larionov acknowledges support from Russian RFBR foundation via grant 09-02-00092. The operation of UMRAO is made possible by funding from the NSF, NASA, and the University of Michigan. The Submillimeter Array is funded by the Smithsonian Institution and the Academia Sinica Institute of Astronomy and Astrophysics. The GASP president acknowledges the ASI support through contract ASI-INAF I/088/06/0."
    },
    "1005/1005.3055_arXiv.txt": {
        "abstract": "{By combining {\\it Herschel}-SPIRE observations obtained as part of the Herschel Virgo Cluster Survey with 21 cm H{\\sc i} data from the literature, we investigate the role of the cluster environment on the dust content of Virgo spiral galaxies. We show for the first time that the extent of the dust disk is significantly reduced in H{\\sc i}-deficient galaxies, following remarkably well the observed `truncation' of the H{\\sc i} disk. The ratio of the submillimetre-to-optical diameter correlates with the H{\\sc i}-deficiency, suggesting that the cluster environment is able to strip dust as well as gas. These results provide important insights not only into the evolution of cluster galaxies but also into the metal enrichment of the intra-cluster medium.} ",
        "introduction": "It is now well established that the evolution of spiral galaxies significantly depends on the environment they inhabit. The reduction in the star formation rate (e.g., \\citealp{lewis02}) and atomic hydrogen (H{\\sc i}) content  (e.g., \\citealp{giova85}) of galaxies when moving from low- to high-density environments indicates that clusters are extremely hostile places for star-forming galaxies. However, a detailed knowledge of the effects of the environment on all the components of the interstellar medium (ISM) is still lacking. Particularly important is our understanding of how the environment is able to affect the dust content of cluster spirals. Dust plays an important role in the process of star formation, since it acts as a catalyzer for the formation of molecular hydrogen (H$_{2}$, from which stars are formed) and prevents its dissociation by the interstellar radiation field. Thus, the stripping of dust might significantly affect the properties of the ISM in infalling cluster spirals.  Since dust is generally associated with the cold gas component of the ISM, it is expected that when the H{\\sc i} is stripped part of the dust will be removed as well, but no definitive evidence of a reduced dust content in cluster galaxies has been found so far. For a fixed morphological type, H{\\sc i}-deficient galaxies\\footnote{The H{\\sc i}-deficiency ($def_{\\rm HI}$) is defined as the difference, in logarithmic units, between the observed H{\\sc i} mass and the value expected from an isolated galaxy with the same morphological type and optical diameter \\citep{haynes}.} appear to have higher IRAS $f(100~\\mu {\\rm m})/f(60~\\mu {\\rm m})$ flux density ratios (i.e., colder dust temperatures, \\citealp{bicay87}) and lower far-infrared (FIR) flux densities per unit optical area \\citep{doyon89} than gas-rich galaxies. However, by using ISO observations of the Virgo cluster \\citep{tuffs02}, \\cite{popescu02} find no strong variation with cluster-centric distance in the dust properties of each morphological type. Only the most extreme H{\\sc i}-deficient galaxies appear to be lacking a cold dust component. More recently, \\cite{review} have revealed an interesting trend of decreasing dust masses per unit of H-band luminosity with decreasing distance from the center of Virgo. Thus, it is still an open issue whether or not dust is removed from infalling cluster spirals.   The launch of {\\it Herschel} \\citep{pilbratt10} has opened a new era in the study of environmental effects on dust. Thanks to its high spatial resolution and sensitivity to all dust components, {\\it Herschel} will be able to determine if cluster galaxies have lost a significant amount of their dust content. Ideally, this analysis should be done on a large, statistically complete sample, following the same criteria used to define the H{\\sc i}-deficiency parameter \\citep{haynes}: i.e., by comparing the dust content of galaxies of the same morphological type but in different environments. By observing a significant fraction ($\\sim$64 deg$^2$) of the Virgo cluster at 100, 160, 250, 350 and 500 $\\mu$m, the Herschel Virgo Cluster Survey (HeViCS, \\citealp{davies10}, hereafter Paper I; see also http://www.hevics.org) will soon provide the optimal sample for such an investigation. In the meantime, with the first HeViCS data it is possible to use a more indirect approach and compare the extent of the dust disk in gas-rich and gas-poor cluster galaxies. Since previous studies have shown that the H{\\sc i} stripping is associated with a `truncation'\\footnote{The term `truncation' is used here to indicate either an abrupt steepening of the surface-brightness profile or, more simply, a significant reduction in the disk scale-length compared to the optical one.} of the gas \\citep{cayatte94} and star-forming disk \\citep{koopmann04,cati05,review}, if the dust follows the atomic hydrogen we should find a reduction in the extent of the dust disk with increasing H{\\sc i}-deficiency. \\begin{figure*} \\centering \\includegraphics[width=17cm]{collage_nolabel.epsi} \\caption{The optical (top), 250 $\\mu$m (middle) and H{\\sc i} (bottom) maps for galaxies in our sample with different degrees of H{\\sc i}-deficiency.} \\label{fc} \\end{figure*}  In this paper we will take advantage of the HeViCS observations, obtained as part of the {\\it Herschel} Science Demonstration (SD) phase, to investigate the correlation between the dust distribution and gas content in cluster galaxies. ",
        "conclusions": "In this paper, we have shown that in H{\\sc i}-deficient galaxies the dust disk is significantly less extended than in gas-rich systems. This result, combined with the evidence that H{\\sc i}-deficient objects show a reduction in their submm-to-K-band flux density ratios, suggests that when the atomic hydrogen is stripped part of the dust is removed as well. However, the dust stripping appears efficient only when very gas-poor spirals are considered, implying that in order to be significant the stripping has to occur well within the optical radius. This is consistent with \\cite{thomas04} who found that the 850 $\\mu$m scale-length of nearby galaxies is smaller than the H{\\sc i}, suggesting that outside the optical radius the gas-to-dust ratio is higher than in the inner parts.  Our analysis provides evidence that the cluster environment is able to significantly alter the dust properties of infalling spirals. We note that this has only been possible thanks to the unique spatial resolution and high sensitivity in detecting cold dust provided by the {\\it Herschel}-SPIRE instrument and to the wide range of H{\\sc i}-deficiencies covered by our sample. Once combined with the direct detection of stripped dust presented by \\cite{cortese10b} and \\cite{gomez10}, our results highlight dust stripping by environmental effects as an important mechanism for injecting dust grains into the intra-cluster medium, thus contributing to its metal enrichment. This is consistent with numerical simulations which predict that ram pressure alone can already contribute $\\sim$10\\% of the enrichment of the ICM in clusters \\citep{domainko}. Interestingly, the stripped grains should survive in the hot ICM long enough to be observed \\citep{popescu00,clemens10}.  Once completed, HeViCS will allow a search for additional evidence of dust stripping and place important constraints on the amount of intra-cluster dust present in Virgo. Moreover, in combination with the Herschel Reference Survey \\citep{hrs}, it will be eventually possible to accurately quantify the degree of dust-deficiency in Virgo spirals."
    },
    "1005/1005.4958_arXiv.txt": {
        "abstract": "A sample consisting of 27 X-ray selected galaxy clusters from the \\textit{XMM-LSS} survey is used to study the evolution in the X-ray surface brightness profiles of the hot intracluster plasma. These systems are mostly groups and poor clusters, with temperatures 0.6-4.8 keV, spanning the redshift range 0.05 to 1.05. Comparing the profiles with a standard \\bmodel\\ motivated by studies of low redshift groups, we find 54\\% of our systems to possess a central excess, which we identify with a cuspy cool core. Fitting \\bmodel\\ profiles, allowing for blurring by the XMM point spread function, we investigate trends with both temperature and redshift in the outer slope ($\\beta$) of the X-ray surface brightness, and in the incidence of cuspy cores. Fits to individual cluster profiles and to profiles stacked in bands of redshift and temperature indicate that the incidence of cuspy cores does not decline at high redshifts, as has been reported in rich clusters. Rather such cores become more prominent with increasing redshift. $\\beta$ shows a positive correlation with both redshift and temperature. Given the $\\beta$-$T$ trend seen in local systems, we assume that temperature is the primary driver for this trend. Our results then demonstrate that this correlation is still present at $z\\gsim0.3$, where most of our clusters reside. ",
        "introduction": "Clusters of galaxies, as the largest virialised gravitationally-bound products of the process of hierarchical structure formation, are powerful probes for both testing cosmological models and tracing structural evolution (e.g., \\citealt{Voit05}). One of the most important properties of galaxy clusters is their mass. Since cluster mass cannot be directly observed, it is studied indirectly through observables such as X-ray radiation emitted by the intracluster medium (ICM) which represents 80\\% of the total baryonic component of galaxy clusters at z=0 (\\citealt{Ettori04}) and accounts for about 10\\% of the total (including dark) mass content of clusters (\\citealt{Sarazin86}). The study of the ICM can provide important insights into the evolution and dynamics of cluster and their member galaxies.  Observationally, there are two distinct classes of clusters: cool core (CC) clusters with dense gaseous core regions in which gas temperature drops inwards, and non-cool core (NCC) clusters with shallower core profiles which often exhibit more internal structure (e.g., \\citealt{Jones84}; \\citealt{Ota04}; \\citealt{Peres98} and \\citealt{Schuecker01}). Cool core clusters have sharply peaked X-ray emission at their centres due to the rise in central gas density which accompanies the central cooling. However, the gas is not observed to cool to very low temperatures at the rates naively expected from the observed core X-ray luminosities, and a consensus has now emerged that this is due to the effects of feedback from a central active galactic nucleus, which limits the effects of cooling through processes which are still not very well understood. For reviews of cool cores in clusters, see for example, \\citealt{Fabian94}, \\citealt{Donahue04} and \\citealt{Peterson06}.  In the local Universe, some studies have found that nearly two thirds of clusters have cool cores (e.g. \\citealt{Peres98}, \\citealt{White97} and \\citealt{Vikhlinin07}). However, other studies gave different values: \\cite{Edge92} found a CC fraction as high as 90\\%, while the results of \\cite{Chen07} indicated that 49\\% of local clusters host cool cores. These differences relate to both the selection of the cluster sample, and the way in which cool cores are identified within them.  The evolution of cooling within clusters provides an important probe of the history of cosmic feedback (\\citealt{Voit05}). At intermediate redshifts ($z\\approx0.15-0.4$), \\cite{Bauer05}, using a sample of 38 X-ray-luminous clusters, found that cool cores appeared still to be common, with an incidence nearly identical to that in luminous low-redshift clusters. Consequently, they suggested that heating and cooling processes must have stabilised in massive clusters since $z\\sim0.4$.  At higher redshifts, \\cite{Vikhlinin07} reported that the fraction of clusters with cuspy X-ray cores dropped from $\\sim$70\\% at $z\\sim0$ to $\\sim$15\\% at $z>0.5$. \\cite{Santos08} compared the fraction of clusters with non-cool cores, moderate cool cores and strong cool cores in nearby ($0.143\\leq z \\leq0.3$) and high redshift ($0.7\\leq z < 1.4$) clusters. These authors detected a significant fraction of clusters harbouring moderate cool cores out to z=1.4, similar to the fraction in their low-redshift sample. However, they noticed an absence of clusters with strong-cool cores at redshift $z>0.7$.   Regarding the spatial distribution of the ICM, \\cite{Cavaliere76} introduced the \\bmodel\\ profile, motivated by the distribution expected for an isothermal plasma in hydrostatic equilibrium with a virialised mass distribution. Although it is now known that the gas is rarely isothermal, the \\bmodel\\ is generally found to give a reasonable representation of X-ray surface brightness profiles (e.g. \\citealt{Neumann99}). However, an additional central component is usually required to fit the inner regions of CC clusters (e.g. \\citealt{Pratt02}), and detailed studies of surface brightness profiles of clusters extending to large radii have shown that the logarithmic slope continues to increase slowly towards larger radii (e.g. \\citealt{Vikhlinin99}, \\citealt{Croston08} and \\citealt{Maughan08}). It is less clear whether this progressive steepening is also present in galaxy groups. For example, \\cite{Rasmussen04} traced the surface brightness out to \\rfh\\ in two rich groups and found them to be well fitted by simple \\bmodel s the whole way.  For \\bmodel\\ fits, the $\\beta$ parameter, which characterises the outer slope of the surface brightness profile, has a value of $\\approx 2/3$ for rich clusters (\\citealt{Jones84}) and lower values for poor clusters and galaxy groups (\\citealt{Finoguenov01}, \\citealt{Helsdon00} and \\citealt{Horner99}). Several studies have shown that $\\beta$ has a mild positive trend with the average temperature of the ICM in nearby clusters (\\citealt{Vikhlinin99}, \\citealt{Croston08} and \\citealt{Maughan08}) and that the value in poorer galaxy groups is lower (i.e. flatter surface brightness slope) than that in clusters (\\citealt{Osmond04b}). In terms of evolution, a study of Chandra data for 115 clusters spanning the range $0.1<z<1.3$ by \\cite{Maughan08} shows some indication that the $\\beta$-$T$ correlation is weaker for clusters at $z>0.5$.  Low-mass galaxy clusters or groups, with ICM temperatures less than 2-3 keV, play an important role in the evolution of galactic systems because they lie at a transition between the field environment and rich cluster environments, and also because non-gravitational processes have a larger impact in groups than in rich clusters (e.g., \\citealt{Zabludoff98}, \\citealt{Ponman03} and \\citealt{Sun09}). However, these poor systems, and in particular the evolution of their properties, have received rather little attention. This is mainly due to the difficulty in detecting and studying them, especially at large redshifts, due to their faint X-ray emission and small complement of galaxies.  For these reasons, research on the evolution of galactic systems in the regime of groups and poor clusters has only started recently, as a result of improvements in observing capabilities in both the X-ray and optical. By comparing optically-selected systems at $0.3 \\leq z \\leq 0.55$ with nearby groups, \\cite{Wilman05} showed that the fraction of group members undergoing significant star formation increases strongly with redshift out to $z\\sim0.45$. However, the study of X-ray selected groups by \\cite{Jeltema07} showed a contrary result: they did not observe significant evolution in the morphology or star formation of the galaxy populations in their $0.2<z<0.6$ groups compared to low-redshift X-ray luminous groups. They argued that this discrepancy could be due to different selection methods, since optically-selected systems are typically lower in mass and contain more spiral galaxies and therefore a stronger evolution in the galaxies is expected. They also found that their moderate redshift groups had galaxy populations similar to clusters at the same redshift; in particular, a large fraction of early-type galaxies and a low fraction of galaxies with significant star formation. However, in contrast to the situation in low redshift X-ray bright groups, a significant fraction of these intermediate redshift groups were found (\\citealt{Mulchaey06}, \\citealt{Jeltema07}) to have no bright early-type galaxy at the centre of the X-ray emission, or to have a central galaxy with multiple nuclei. This was taken as evidence for the dynamical youth of many of these groups.  The small number of studies which have addressed the evolution in the X-ray properties of galaxy groups have found little convincing evidence for any. \\cite{Jeltema06}, in a multiwavelength study of six galaxy groups and poor clusters at intermediate redshift (0.2-0.6), found that they appear to follow the scaling relations between luminosity, temperature, and velocity dispersion defined by low-redshift groups and clusters. This is also true (\\citealt{Jeltema09}) for three higher redshift poor clusters from the AEGIS survey. A study of evolution in the $L$-$T$ relation based on the present XMM-LSS cluster sample by \\cite{Pacaud07}, taking into account the selection function of the survey, found that the range of models consistent with the data included self-similar evolution, and also (marginally) a no-evolution model. \\cite{Finoguenov07} extracted a larger sample of 72 clusters (mostly poor ones) from the XMM-Newton observations of the COSMOS field, and found no evidence for evolution in the luminosity function of these systems out to $z\\sim 1$, though the quality of their data did not permit them to study the morphology of the X-ray emission.  Motivated by the the paucity of information available for the evolution of the ICM in the important environment of low-mass galaxy clusters, we aim in this study to shed light on the spatial distribution of the ICM in X-ray selected clusters covering a wide redshift range ($z\\sim 0-1$), paying special attention to trends in the slope and central cuspiness of the X-ray emission.  The paper is constructed as follows: in section 3.2, we describe the data and briefly introduce the properties of the cluster sample; then we describe the data reduction used to produce X-ray surface brightness profiles. In section 3.3, we present our results, starting with the individual cluster profiles, and then profiles of redshift-stacked and temperature-stacked clusters. In section 3.4, we discuss the implication of our results and compare them with other studies. Finally, in section 3.5, we summarise our conclusions.  Throughout this paper, we adopt  the cosmological parameters from the five-year WMAP data presented by \\cite{Hinshaw09}, namely: $H_0=70.5$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m=0.27$, $\\Omega_\\Lambda=0.73$. ",
        "conclusions": "In this paper, we used \\textit{XMM-Newton} observations of 27 X-ray selected galaxy clusters spanning the redshift range ($0.05 \\leq z \\leq 1.05$) to study the spatial properties of their ICM. Most of these clusters fall in the realm of low-mass clusters or groups, with ICM temperatures from 0.63 to 4.80~keV. The XMM data provide typically a few hundreds X-ray source counts. We extracted and vignetting-corrected the profiles to 3$\\times$\\rfh\\ where they flattened and reached the photon background values, which were estimated locally for each cluster.  In addition to the individual profiles, we also stacked the profiles into three redshift and temperature bins. To explore the effects of \\textit{Malmquist bias}, we further stacked clusters with similar redshifts/temperatures into two subsets each with different averaged temperature/redshift. Both individual and stacked profiles were fitted with blurred (to account for the PSF errors of the \\textit{XMM-Newton} cameras) \\bmodel s with both free and fixed core radii. The fixed-\\rc\\ \\bmodel\\ fits were used to test whether a profile showed evidence of a cuspy core, making this study the first to probe the evolution of CCs out to $z\\gsim0.3$ within poor clusters.  Our main conclusions are: \\begin{itemize} \\item We find that 54\\% of our sample show evidence for cool cores, in the form of a central excess (at $>1\\sigma$ significance) above a standard \\bmodel. \\item For the free-\\rc\\ fits to individual clusters, the median value of $\\beta$ is 0.61, and the median \\rc\\ is 0.08$\\times$\\rfh. \\item For the fixed-\\rc\\ fits to individual clusters, the median $\\beta$ is 0.63. \\item Twelve systems in our sample (with $\\overline{kT}=1.69$ keV) have $z\\sim0.3$, allowing us to characterise the X-ray surface brightness profiles of intermediate redshift X-ray selected groups.  The free-fit parameters to the stacked data from these 12 systems gives $\\beta=0.51\\pm0.01$ and \\rc=$0.06\\pm0.01\\times$\\rfh. This stacked profile indicates the presence of CCs ($f_c=1.56\\pm0.11$), with 7 of the 12 systems showing a significant central excess in their individual profiles. \\item Stacked and individual profiles for our sample of poor galaxy clusters show that the CCs do not disappear at high redshift, but rather become more prominent, though one would like to confirm this result with higher spatial resolution observations. \\item The slope parameter, $\\beta$, shows a positive trend with both redshift and temperature in our data. Combining this results with previous findings, we incline towards a trend with temperature (and hence mass) rather than redshift. The present study then demonstrates for the first time, that the $\\beta$-$T$ trend seen at low $z$ is also present in groups and clusters at $z\\gsim0.3$. \\end{itemize}"
    },
    "1005/1005.0056_arXiv.txt": {
        "abstract": "We report the Suzaku observation of 1E~1740.7$-$2942, a black hole candidate called the ``Great Annihilator'' (GA). The high-quality spectrum of Suzaku provides the severest constraints on the parameters of the GA. Two clumpy structures are found around the GA in the line images of Fe\\emissiontype{I} K$\\alpha$ at 6.4~keV and S\\emissiontype{XV} K$\\alpha$ at 2.45~keV. One clump named M\\,359.23$-$0.04 exhibits the 6.4-keV line with an equivalent width of $\\sim$ 1.2~keV, and is associated with a molecular cloud in the radio CS($J$=1--0) map. Thus the 6.4-keV line from M\\,359.23$-$0.04 is likely due to X-ray fluorescence irradiated by an external X-ray source. The irradiating X-rays would be either the past flare of Sagittarius\\,A$^*$ or the bright nearby source, the GA. The other clump named G\\,359.12$-$0.05 is associated with the radio supernova remnant candidate G\\,359.07$-$0.02. We therefore propose that G\\,359.12$-$0.05 is an X-ray counterpart of G\\,359.07$-$0.02. G\\,359.12$-$0.05 has a thin thermal plasma spectrum with a temperature of $kT \\sim 0.9$~keV. The plasma parameters of G\\,359.12$-$0.05 are consistent with those of a single supernova remnant in the Galactic center region. ",
        "introduction": "The Galactic center (GC) region has many celestial objects such as a supermassive black hole Sagittarius\\,A$^*$ (Sgr\\,A$^*$), star-forming regions, dense molecular clouds (MCs). The GC region is also full of high-energy phenomena. For example, several supernova remnants (SNRs) or candidates have been recently found (e.g., \\cite{Nobukawa.2008}; \\cite{Mori.2008}; \\cite{Sawada.2009}; \\cite{Tsuru.2009}). The most characteristic phenomenon is the diffuse 6.7-keV lines from highly ionized irons (\\cite{Koyama.1989}; \\cite{Yamauchi.1990}). The origin of this emission is probably due to a plasma with $kT~\\sim$~6.5 keV temperature \\citep{Koyama.2007b}, but is still a debatable issue. Another notable feature in the GC region is the clumpy 6.4-keV lines from neutral irons (e.g., \\cite{Koyama.1996}). The origin is considered as MCs illuminated by X-rays or electrons. Some of the 6.4-keV clumps are well explained by the irradiation of the past X-ray flare from Sgr\\,A$^*$ (e.g., \\cite{Nobukawa.2008}; \\cite{Inui.2009}; \\cite{Nakajima.2009}).  1E~1740.7$-$2942 was discovered by the Einstein observatory \\citep{Hertz.1984} in the direction of the GC region, and was found to be the brightest GC source at above 20~keV (\\cite{Skinner.1987}; \\cite{Sunyaev.1991}). The time variation of the flux and the spectrum of 1E~1740.7$-$2942 are similar to those of Cygnus\\,X-1, the archetypal black hole candidate, and hence 1E~1740.7$-$2942 has been considered to be a black hole candidate (\\cite{Cook.1991}; \\cite{Skinner.1991}).  GRANAT detected a prominent bump on the spectrum of 1E~1740.7$-$2942 at 300--600 keV, which was interpreted as an electron-positron annihilation line at 511~keV (\\cite{Bouchet.1991}; \\cite{Sunyaev.1991}; \\cite{Churazov.1993a}; \\cite{Cordier.1993}). Thus 1E~1740.7$-$2942 is named the ``Great Annihilator'' (GA). However, no evidence for the annihilation line has been found so far by other satellites (\\cite{Harris.1994}; \\cite{Jung.1995}; \\cite{Smith.1996}; \\cite{Cheng.1998}; \\cite{Bouchet.2009}).  The VLA radio observations discovered a radio counterpart of the GA and non-thermal double jet-like structures emanating from the GA \\citep{Mirabel.1992}, and hence the GA is a ``micro quasar''. Like the other Galactic jet sources, the GA would be a binary system with a stellar-mass black hole. In fact, the possible orbital period of 12.73 days was discovered by \\citet{Smith.2002}. However, no clear companion star was found with the optical, infrared or radio observations, due mainly to the strong interstellar absorption toward the GC region (e.g., \\cite{Eikenberry.2001}).  The X-ray spectrum of the GA below 10 keV is explained by an absorbed power-law model (\\cite{Sakano.1999}; \\cite{Gallo.2002}). \\citet{Cui.2001} reported the extended X-ray emission which is perpendicular to the radio-jets (but see \\cite{Gallo.2002}). On the other hand, the INTEGRAL observations suggested that the spectrum in the 10--100~keV band is explained by a model of either a comptonized-plasma plus a power-law or two comptonized-plasmas \\citep{Bouchet.2009}.  The radio observations found a giant MC near the GA (\\cite{Bally.1991}; \\cite{Mirabel.1991}). Then the authors proposed the Bondi-Hoyle accretion mechanism \\citep{Bondi.1944}; the GA is in the MC, and is powered by the gas accretion from the MC. However the ASCA observations found no clear evidence for dense gas around the GA (\\cite{Churazov.1996}; \\cite{Sakano.1999}).  In order to investigate the accretion mechanisms, whether the Bondi-Hoyle process, due to a binary companion or else, we performed the Suzaku observation on the GA, with its high-quality spectrum and low background.  The other objective of the observation is to  discover local structures, if any, on the GA and in the close vicinity of the GA, and investigate the physical relation to the GA. We assume the distance to the GC to be 8.5~kpc in this paper. ",
        "conclusions": "\\subsection{Nature of M\\,359.23$-$0.04}  From the radio MC observation of CS(\\textit{J}=1--0) \\citep{Tsuboi.1999}, we find that at least one clump of the MCs is associated with M\\,359.23$-$0.04 in the two velocity bands: the $-$140~to~$-$120~km~s$^{-1}$ band and $-$20~to~0~km~s$^{-1}$ band. Figure \\ref{figure7} shows the radio contours overlaid on the X-ray images. We find several clumps of MCs and designate them as MC1--6. Among them, MC2 and MC3 are possibly associated with M\\,359.23$-$0.04.  The best-fit \\textit{EW}$_{\\mathrm{6.4\\,keV}}$ of M\\,359.23$-$0.04 is 1.2--1.4~keV (see table \\ref{table4} and section 3.5). This value is naturally explained by the X-ray irradiation (the X-ray reflection nebula: XRN model) on the MC with one solar abundance (e.g., \\cite{Murakami.2000}; \\cite{Nakajima.2009}). If we adopt the electron bombarding model (\\cite{Yusef-Zadeh.2007}; \\cite{Fukuoka.2009}), we need more than 3 solar abundances. The infrared observations (e.g., \\cite{Davies.2009}) argued that the iron abundance of the GC is nearly equal to the solar value. In fact, most of the Suzaku observations of the GC region reported the abundances between 1--2 solar (e.g., \\cite{Mori.2008}, \\cite{Nobukawa.2008}; \\cite{Tsuru.2009}; \\cite{Nobukawa.2010}). Our observation also implies  1.2--1.5 solar abundances in this region (see table \\ref{table2} and \\ref{table3}). The XRN model is therefore more favorable for the M\\,359.23$-$0.04 X-ray emission.  \\begin{figure}[t] \\begin{center} \\FigureFile(80mm,120mm){Fig7.eps} \\end{center} \\caption{The 6.4-keV line images (same as figure \\ref{figure3}b) are shown in the gray scale. The blue contours trace the CS(\\textit{J}=1-0) radio observation \\citep{Tsuboi.1999} in the velocity bands of $-$140~to~$-$120~km~s$^{-1}$ (a), and in the $-$20~to~$-$0~km~s$^{-1}$ (b). The green ellipse shows the source region of M\\,359.23$-$0.04. The red circle and cross indicate the \\timeform{3'}-radius region and the position of the GA, respectively. Individual prominent spots (clumps) in the radio contours are labeled as MC1 to MC6 (see text).}\\label{figure7} \\end{figure}  In the XRN model, the irradiating source should be very bright or be located near at M\\,359.23$-$0.04. Thus the GA at a projected distance of \\timeform{6'} from M\\,359.23$-$0.04 and the past bright Sgr\\,A$^*$ (e.g., \\cite{Inui.2009}) are the possible candidates for the irradiating source. Another possible candidate is a transient source. Although this possibility due to an unknown transient is not excluded, the available observations place this possibility to be a less plausible hypothesis, because no bright transient source has been so far reported from the close vicinity of M\\,359.23$-$0.04 (e.g.,\\cite{in't Zand.2001}; \\cite{Sakano.2002}; \\cite{Muno.2009}).  Based on the observed profile (figure \\ref{figure3}b), the projected shape of M\\,359.23$-$0.04 is approximated with an ellipse of $12.5~\\mathrm{pc} \\times 22.5~\\mathrm{pc}$ (minor $\\times$ major axis). Adopting a reasonable assumption that the 3-dimensional structure of M\\,359.23$-$0.04 is a prolate spheroid with the line-of-sight depth of 12.5~pc (the same as the minor axis), we estimate the volume to be $5 \\times 10^{58}~\\mathrm{cm}^{3}$ (a prolate spheroid with axises of $12.5~\\mathrm{pc} \\times 12.5~\\mathrm{pc} \\times 22.5~\\mathrm{pc}$). The observed \\textit{N}$_{\\mathrm{H}}$ is $(19\\textrm{--}35)~\\times 10^{22}$~cm$^{-2}$. On the other hand, the past studies reported that the interstellar absorption to the GC region is \\textit{N}$_{\\mathrm{H}} = 6~\\times 10^{22}$~cm$^{-2}$ (\\cite{Sakano.2002}; \\cite{Ryu.2009}). To confirm this value, we conduct the model fitting of the GCDX in the close vicinity of the GA. We then find that the interstellar absorption is \\textit{N}$_{\\mathrm{H}} = 5.7$~(5.5--6.0)~$\\times 10^{22}$~cm$^{-2}$, which is consistent with the values in the references. Subtracting the interstellar absorption to the GC region, we estimate that the intrinsic \\textit{N}$_{\\mathrm{H}}$ is about $(1\\textrm{--}3) \\times 10^{23}$~cm$^{-2}$. On the other hand, \\textit{N}$_{\\mathrm{H}}$ of MC2 or MC3, a counterpart of M\\,359.23$-$0.04 is $\\sim 1 \\times 10^{23}~$cm$^{-2}$ from the CS(\\textit{J}=1--0) observation \\citep{Tsuboi.1999}. Thus \\textit{N}$_\\mathrm{H}$ of the cloud is in the range of (1--3) $\\times 10^{23}$~cm$^{-2}$.  Adopting these physical parameters and assuming that the distances of M\\,359.23$-$0.04 from the GA and Sgr\\,A$^*$ are the same as the projected distances, we estimate the required luminosity of the GA and Sgr\\,A$^*$ to produce the 6.4-keV-line flux of M\\,359.23$-$0.04. In this process, the spectrum shapes of the GA and Sgr\\,A$^*$ are assumed to be a power-law with photon index  $\\Gamma$ = 1.4 and 2, respectively (see table \\ref{table2} and \\cite{Murakami.2000}). The 6.4-keV-fluorescence line is due to X-rays above the iron K-edge energy at 7.1~keV \\citep{Murakami.2000}. At these high energy X-rays, we can ignore the interstellar absorption between the irradiating source and the XRN. The results of the required luminosity in the 2.5--12~keV band for the GA and Sgr\\,A$^*$ are $4 \\times 10^{36}~\\mathrm{erg~s}^{-1}$ and $5 \\times 10^{38}~\\mathrm{erg~s}^{-1}$, respectively.  The observed luminosity of the GA is smaller by a factor of 2 than the required luminosity (see table \\ref{table2}). This difference is within the systematic errors of the size, density and distance; i.e., the real distance would be larger than the projected distance. In addition, the flux-variation of the GA  by a factor of 2--5 was reported in the long time period (\\cite{Churazov.1993b}; \\cite{Sakano.1999}; \\cite{Smith.2002}). Thus the scenario of the GA origin is consistent with the observed results.  The other possibility, the past flare of Sgr\\,A$^*$,  has been proposed for many other 6.4-keV clumps in the GC region (e.g., \\cite{Nobukawa.2008}; \\cite{Nakajima.2009}). The same argument can also be applied to M\\,359.23$-$0.04, and thus the scenario of Sgr\\,A$^*$ origin is possible.  In figure \\ref{figure7}, we find other 6.4-keV clumps at near MC1 and MC4. The 6.4-keV fluxes of them are $\\sim$ 2 times fainter than that of M\\,359.23$-$0.04, and hence the physical association is less clear. The other MCs, MC5 and MC6, show no association to the 6.4-keV emissions. These features may not be surprising, if the putative past flare was variable in the time scale of a few years \\citep{Koyama.2009}.  On the other hand, the scenario of the GA origin can be more easily assessed, if the real distances (not projected distances) of these clumps from the GA are 1.4 times larger than those of MC2 or MC3, because the 6.4-keV fluxes from  MC1, MC4, MC5, and MC6 are less than half of those from MC2 or MC3.    \\subsection{Nature of the GA}  The observed column density of the GA is $\\textit{N}_{\\mathrm{H}} = 12 \\times 10^{22}~\\mathrm{cm}^{-2}$. If the GA is associated with M\\,359.23$-$0.04 (see section 4.1), the GA lies in the GC region, because \\citet{Tsuboi.1999} showed that MCs of CS(\\textit{J}=1--0) are in the GC region. In this case, the large $\\textit{N}_{\\mathrm{H}}$ of the GA is explained by intrinsic absorption \\citep{Kawai.1988} of the order of $10^{23}$~cm$^{-2}$.  We constrain the upper-limit for the \\textit{EW}$_\\mathrm{6.4\\,keV}$ of 8~eV. The iron abundance is also accurately determined to be 1.38--1.44 solar. If the GA is fully covered by the gas for $\\Omega$ = 4$\\pi$ steradian, the upper-limit of \\textit{EW}$_\\mathrm{6.4\\,keV} = 8~\\mathrm{eV}$ constrains the column density of the circumstellar gas within 7.5~pc (\\timeform{3'} at 8.5 kpc) to be \\textit{N}$_{\\mathrm{H}}$ $\\leq$ 4$\\times$ 10$^{21}$~cm$^{-2}$ (\\cite{Inoue.1985}). The implication of the small value of \\textit{EW}$_\\mathrm{6.4\\,keV}$ was discussed in detail by  \\citet{Sakano.1999}.  Here we discuss briefly following \\citet{Sakano.1999}, based on the more severe constraint of \\textit{EW}$_\\mathrm{6.4\\,keV}$ (\\textit{EW}$_\\mathrm{6.4\\,keV}$ = 15 eV in \\cite{Sakano.1999} \\textit{vs} 8 eV in our case). The \\textit{EW}$_\\mathrm{6.4\\,keV}$ of 8 eV implies that the local density near the GA is less than 2 $\\times$ 10$^{2}$~H~cm$^{-3}$. From figure 6 of \\citet{Sakano.1999}, the velocity of the GA must be less than 7~km~s$^{-1}$ under the Bondi-Hoyle accretion mechanism even if the GA has the mass of 20~\\Mo. Since no data of the velocity dispersion for a black hole are available, we refer to that of the radio pulsars (neutron stars) from \\citet{Faucher.2006}. We then conclude that the possibility of the GA with a smaller velocity than 7~km~s$^{-1}$ is only $2 \\times 10^{-3}$\\%. Thus the Bondi-Hoyle accretion for the GA is extremely unlikely.  Although  the observed \\textit{EW}$_\\mathrm{6.4\\,keV} < 8~\\mathrm{eV}$ gives the column density of the circumstellar gas within 7.5~pc to be \\textit{N}$_{\\mathrm{H}}$ $\\leq$ 4$\\times$ 10$^{21}$~cm$^{-2}$ (this value is averaged over 4$\\pi$ steradian), the intrinsic \\textit{N}$_\\mathrm{H}$ determined from the low energy absorption is in the order of $10^{23}$~cm$^{-2}$. Thus the large circumstellar gas should be localized in front of the GA. \\citet{Bally.1991} reported that an MC is associated to the GA, but no MC is found in the line of sight to the GA from \\citet{Tsuboi.1999}. This inconsistency may come from the poor spatial resolution of the former observation. In fact, the spatial resolution of the former observation is $\\sim$~\\timeform{2'}, while the latter is $\\sim$~\\timeform{0.5'}. Our implication is consistent with \\citet{Tsuboi.1999}, and hence one possible scenario for the large absorption is due to an accreting gas; the GA is a binary system with a nearly edge-on accretion disk.  \\subsection{Plasma Parameters of G\\,359.12$-$0.05}  \\begin{figure}[t] \\begin{center} \\FigureFile(80mm,65mm){Fig8.eps} \\end{center} \\caption{The 2.45-keV line image (same as figure \\ref{figure3}a) is shown in the gray scale. The magenta contours trace G\\,359.07$-$0.02 from the 90~cm radio continuum observation \\citep{LaRosa.2000}. The green ellipse shows the source region of G\\,359.12$-$0.05. The red circle and cross indicate the \\timeform{3'}-radius region and the position of the GA, respectively.}\\label{figure8} \\end{figure}  The observed column density of G\\,359.12$-$0.05 $\\textit{N}_{\\rm{H}} = 6.7 \\times 10^{22}~\\mathrm{cm}^{-2}$ is the same as the GC typical column density of $N_{\\rm{H}} = 6 \\times 10^{22}~\\mathrm{cm}^{-2}$, and hence G\\,359.12$-$0.05 is likely to be in the GC region. Since the source area is $\\sim$ $140~\\mathrm{arcmin}^2 = 8.1 \\times 10^{39}~\\mathrm{cm}^{2}$ (at 8.5~kpc), the volume of the plasma is estimated as $\\sim$ $(8.1~\\times~10^{39}~\\mathrm{cm}^2)^{\\frac{3}{2}} = 7.3 \\times 10^{59}~\\mathrm{cm}^{3}$. Using the filling factor of $f$, we estimate the electron density, total mass, and total thermal energy to be  $n_\\mathrm{e}$ $\\sim$ 0.3~$f^{-\\frac{1}{2}}$~cm$^{-3}$, $M \\sim$ 2 $\\times 10^{2}$~$f^{\\frac{1}{2}}$~\\Mo, and $E_{\\mathrm{th}}$ $\\sim$ 1 $\\times$ 10$^{51}$~$f^{\\frac{1}{2}}$~erg, respectively.  The estimated mass and thermal energy are slightly larger than those of the other SNRs in the GC region, but are consistent with those of a single SNR. As shown in figure \\ref{figure8}, the radio SNR candidate G\\,359.07$-$0.02 \\citep{LaRosa.2000} coincides in position with G\\,359.12$-$0.05. We therefore propose that G\\,359.12$-$0.05 is an X-ray counterpart of G\\,359.07$-$0.02.  The thin thermal plasma spectrum and the physical parameters place G\\,359.12$-$0.05 to be an X-ray SNR in the GC region. From the size of $\\sim$ $9.0~\\times~10^{19}~\\mathrm{cm}$ and the sound velocity of $\\sim$ $5~\\times~10^{7}~\\mathrm{cm}~\\mathrm{s}^{-1}$ ($kT \\sim0.9$ keV), the age of this SNR is estimated to be $\\sim$ $6~\\times~10^{4}$ years.  If the GA is also located in the GC region (see section 4.2), G\\,359.12$-$0.05 may be physically associated with the GA. Then the GA would be a rare system which shows the association of a black hole and an SNR. A similar system is SS433 and the radio SNR W50 (e.g., \\cite{Safi-Harb.1997}).  The nature of SS433 is debatable but is a potential candidate of a black hole. On the other hand, the GA is a more established black-hole candidate. Thus our result provides a unique case of the association between a black hole and an X-ray SNR."
    },
    "1005/1005.5390_arXiv.txt": {
        "abstract": "Numerical models of the tidal disruption of the Sagittarius (Sgr) dwarf galaxy have recently been developed that for the first time simultaneously satisfy most observational constraints on the angular position, distance, and radial velocity trends of both leading and trailing tidal streams emanating from the dwarf. We use these dynamical models in combination with extant 3-D position and velocity data for Galactic globular clusters and dSph galaxies to identify those Milky Way satellites that are likely to have originally formed in the gravitational potential well of the Sgr dwarf, and have been stripped from Sgr during its extended interaction with the Milky Way. We conclude that the globular clusters Arp 2, M 54, NGC 5634, Terzan 8, and Whiting 1 are almost certainly associated with the Sgr dwarf, and that Berkeley 29, NGC 5053, Pal 12, and Terzan 7 are likely to be as well (albeit at lower confidence). The initial Sgr system therefore may have contained 5-9 globular clusters, corresponding to a specific frequency $S_N = 5 - 9$ for an initial Sgr luminosity $M_V = -15.0$. Our result is consistent with the $8\\pm2$ genuine Sgr globular clusters expected on the basis of statistical modeling of the Galactic globular cluster distribution and the corresponding false-association rate due to chance alignments with the Sgr streams. The globular clusters identified as most likely to be associated with Sgr are consistent with previous reconstructions of the Sgr age-metallicity relation, and show no evidence for a second-parameter effect shaping their horizontal branch morphologies. We find no statistically significant evidence to suggest that any of the recently discovered population of ultra-faint dwarf galaxies are associated with the Sgr tidal streams, but are unable to rule out this possibility conclusively for all systems. ",
        "introduction": "Globular clusters tend to form in all sufficiently massive galactic systems during episodes of major star formation, their relatively simple stellar populations providing important constraints on starburst and galaxy formation models (see, e.g., Brodie \\& Strader 2006). While no globular clusters are found in low mass ($M_V \\gtrsim -12$) dwarf galaxies, the more massive dwarfs such as the LMC ($M_V = -18.5$), SMC ($M_V = -17.1$) and Fornax ($M_V = -13.1$) are observed to contain nineteen, eight, and five  globular clusters respectively (see, e.g., Forbes et al. 2000). It is generally accepted that the Sagittarius (Sgr) dwarf spheroidal galaxy ($M_V = -13.64$; Law \\& Majewski 2010 [hereafter LM10]) contains at least four clusters (M 54, Arp 2, Terzan 7, Terzan 8; Da Costa \\& Armandroff 1995) within its main body. Since Sgr is known to have experienced significant tidal mass loss, however, it is likely that additional globular clusters may have been stripped from Sgr during its prolonged interaction with the Milky Way and now lie scattered throughout the Galactic halo.  In the anticipation that such clusters are likely distributed across the sky, various authors (e.g., Lynden-Bell \\& Lynden-Bell 1995; Irwin 1999; Bellazzini et al. 2002; Newberg et al. 2003; Majewski et al. 2004; Carraro et al. 2007) have suggested a number of additional globular clusters (e.g., Pal 2, Pal 12, NGC 5634, NGC 2419, Whiting 1, etc.) that may have formed in the Sgr dSph (see Table \\ref{prevcands.table} for a summary). Perhaps the most comprehensive efforts to conduct a systematic census to date have been undertaken by Palma et al. (2002) and Bellazzini et al. (2003a), who demonstrated that there is a strong correlation between the positions and radial velocities of outer halo (Galactocentric radius $r_{\\rm GC} > 10$ kpc) globular clusters and the orbit of Sgr. Bellazzini et al. (2003a) in particular presented a statistical assessment of the significance of the total number of Galactic globular clusters aligned with the Sgr orbital plane, concluding that there was a  $\\lesssim 2$\\% probability that the observed grouping of clusters about the Sgr plane arises by chance. While this approach was certainly effective, it was based on the projected {\\it orbit} of Sgr (as derived by Ibata \\& Lewis 1998) and therefore did not adopt an optimal membership criterion because tidal debris does not precisely follow the orbit of the parent object: Trailing tidal arms lie outside the orbital path of the parent while leading tidal arms lie interior to the orbit and, in the case of Sgr, wrap up around the Galactic Center (see discussion by Johnston et al. 1995, 1999; Choi et al. 2007; Eyre \\& Binney 2009). Moreover, Bellazini et al. (2003a) concentrated on the past orbit of Sgr, an approach that misses potential connections to the leading arm tidal debris.  In this contribution we wish to update and extend these previous analyses, and ask whether each individual Milky Way satellite matches the angular position, distance, and radial velocity of Sgr tidal debris sufficiently well that it is statistically likely to be physically associated with the stream and might therefore have originated in the Sgr dwarf. Such an undertaking is greatly aided by the significantly improved picture of the Milky Way --- Sgr system that has been provided in recent years by the Two Micron All-Sky Survey (2MASS) and Sloan Digital Sky Survey (SDSS) which have mapped the stellar streams from the Sgr dwarf wrapping fully $360^{\\circ}$ across the sky (Majewski et al. 2003; Belokurov et al. 2006a; Yanny et al. 2009).  These surveys have also revealed a wealth of additional substructure in the Galactic halo in the form of both globular clusters and a population of dwarf galaxies with masses $\\sim 10^7 M_{\\odot}$ within their central 300 pc (Strigari et al. 2008) but ultra-faint luminosities comparable to those of globular clusters. The origin of these ultra-faint dwarfs is unknown, and we consider whether any of them may be dynamically associated with the Sgr dSph.  Our effort is also aided by the recent construction of a numerical model of the Sgr stream that reproduces the majority of the observed characteristics of the stellar tidal streams with extremely high fidelity (Law et al. 2009; LM10). The comprehensive view of both leading and trailing arm material (which overlap each other along the line of sight for large swathes of the sky) provided by this model allows us to view potential Sgr members in the context of the full debris system, and assess possible associations with leading and trailing arms based not only on spatial distributions, but on dynamical properties such as radial velocities and proper motions as well.   \\begin{deluxetable*}{llccccccccc} \\tablecolumns{11} \\tablewidth{0pc} \\tabletypesize{\\scriptsize} \\tablecaption{Candidate Sgr Stream Globular Clusters Proposed in the Literature\\tablenotemark{a}} \\tablehead{\\colhead{ID} & \\colhead{Name} & \\colhead{DA95} & \\colhead{LL95} & \\colhead{I99} & \\colhead{D01} & \\colhead{P02} & \\colhead{B03} & \\colhead{M04} & \\colhead{C07} & \\colhead{LM10b}} \\startdata 45\t&\tArp 2\t\t& YES\t& YES\t& YES\t& ...\t\t& YES\t& YES\t& YES\t& ...\t\t& YES\\\\% 52\t&\tM 2\t\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& YES\t& ...\t\t& ...\t\t& NO\\\\ ...\t&\tM 5\t\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& YES\t& ...\t\t& ...\t\t& ...\t\t& NO\\\\ 21\t&\tM 53\t\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& YES\t& ...\t\t& ...\t\t& ...\t\t& NO\\\\ 42\t&\tM 54\t\t\t& YES\t& YES\t& YES\t& ...\t\t& YES\t& YES\t& YES\t& ...\t\t& YES\\\\% 1\t&\tNGC 288\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& YES\t& ...\t\t& ...\t\t& NO\\\\ 13\t&\tNGC 2419\t& ...\t\t& ...\t\t& YES\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& NO\\\\ 18\t&\tNGC 4147\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& NO\t& YES\t& ...\t\t& ...\t\t& NO\\tablenotemark{b}\\\\% 22\t&\tNGC 5053\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& NO\t& YES\t& ...\t\t& ...\t\t& YES\\\\% 26\t&\tNGC 5466\t& ...\t\t& ...\t\t& ...\t\t& NO\t& YES\t& YES\t& ...\t\t& ...\t\t& NO\\\\ 27\t&\tNGC 5634\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& YES\t& ...\t\t& ...\t\t& YES\\\\% 30\t&\tNGC 5824\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& YES\t& ...\t\t& ...\t\t& NO\\\\ 40\t&\tNGC 6426\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& YES\t& ...\t\t& ...\t\t& NO\\\\ ...\t&\tNGC 6356\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& YES\t& ...\t\t& ...\t\t& ...\t\t& NO\\\\ 7\t&\tPal 2\t\t\t& ...\t\t& YES\t& YES\t& ...\t\t& ...\t\t& YES\t& YES\t& ...\t\t& NO\\tablenotemark{b}\\\\% 31\t&\tPal 5\t\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& NO\t& YES\t& ...\t\t& ...\t\t& NO\\\\ 53\t&\tPal 12\t\t& ...\t\t& ...\t\t& YES\t& ...\t\t& YES\t& YES\t& YES\t& ...\t\t& YES\\\\% 54\t&\tPal 13\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& NO\t& YES\t& ...\t\t& ...\t\t& NO\\\\ 19\t&\tRup 106\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& YES\t& ...\t\t& ...\t\t& NO\\\\ ...\t&\tTerzan 3\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& YES\t& ...\t\t& ...\t\t& NO\\\\ 44\t&\tTerzan 7\t\t& YES\t& YES\t& YES\t& ...\t\t& YES\t& YES\t& YES\t& ...\t\t& YES\\\\% 46\t&\tTerzan 8\t\t& YES\t& NO\t& YES\t& ...\t\t& YES\t& YES\t& YES\t& ...\t\t& YES\\\\% 3\t&\tWhiting 1\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& ...\t\t& YES \t& YES\\\\% \\enddata \\tablenotetext{a}{Columns represent Da Costa \\& Armandroff (1995; DA95), Lynden-Bell \\& Lynden-Bell (1995; L95), Irwin et al. (1999; I99), Dinescu et al. (2001; D01), Palma et al. (2002; P02), Bellazzini et al. (2003a; B03), Majewski et al. (2004; M04), Carraro et al. (2007; C07), this contribution (LM10b).  `YES' indicates possible candidate, `NO' indicates disfavored by indicated study.} \\tablenotetext{b}{As discussed in \\S \\ref{discussion.sec}, NGC 4147 and Pal 2 may plausibly be associated with the Sgr stream, but at low statistical confidence.} \\label{prevcands.table} \\end{deluxetable*}     This paper is organized as follows.  In \\S \\ref{model.sec} we describe the general characteristics of the Sgr tidal streams and the $N$-body model that best reproduces its observed properties.  In \\S \\ref{gcsel.sec} we describe our sample of Galactic satellites, incorporating globular clusters, ultra-faint dwarfs, and two open clusters for which association with Sgr has previously been claimed in the literature. In \\S \\ref{quantifying.sec} we describe our numerical technique for quantifying the association of a Galactic satellite with the Sgr stream, applying this to individual stellar clusters and ultra-faint dwarfs in \\S \\ref{compareGC.sec} and \\S \\ref{compareUF.sec} respectively. We discuss the resulting implications for the star cluster budget of Sgr and its contributions to the Galactic halo in \\S \\ref{discussion.sec}, summarizing our conclusions in \\S \\ref{summary.sec}.  We adopt the heliocentric Sgr coordinate system ($\\Lambda_{\\odot}, B_{\\odot}$) defined by Majewski et al. (2003) in which the longitudinal coordinate $\\Lambda_{\\odot} = 0^{\\circ}$ in the direction of Sgr and increases along trailing tidal debris, and $B_{\\odot}$ is positive towards the orbital pole $(l,b)_{\\rm pole} \\approx (274^{\\circ}, -14^{\\circ})$ (see discussion by LM10).  All radial velocities are given in the Galactic Standard of Rest (GSR) frame, with respect to which the Sun has a peculiar motion ($U,V,W$) = ($9,12+220,7$) km s$^{-1}$ and is located $R_{\\odot} = 8$ kpc from the Galactic Center. ",
        "conclusions": "\\label{discussion.sec}   \\begin{deluxetable}{rlc} \\tablecolumns{3} \\tablewidth{0pc} \\tabletypesize{\\scriptsize} \\tablecaption{Star Clusters Associated with the Sgr Stream} \\tablehead{\\colhead{ID} & \\colhead{Name} & \\colhead{Arm} } \\startdata \\multicolumn{3}{c}{High Confidence}\\\\ \\hline 3\t & \tWhiting 1\t & \tT1\\\\ 27\t & \tNGC 5634\t & \tT1\\\\ 42\t & \tM 54\t & \tCORE\\\\ 45\t & \tArp 2\t & \tT1\\\\ 46\t & \tTerzan 8\t & \tT1\\\\ \\hline \\multicolumn{3}{c}{Moderate Confidence}\\\\ \\hline 11\t& \tBerkeley 29\t& L1\\\\ 22\t & \tNGC 5053\t & \tT1\\\\ 44\t & \tTerzan 7\t & \tT1\\\\ 53\t & \tPal 12\t & \tT1/L1\\\\ \\hline \\multicolumn{3}{c}{Low Confidence}\\\\ \\hline 7\t&\tPal 2\t\t& T1\\\\ 18\t&\tNGC 4147\t& T1\\\\ \\enddata \\label{gcbudget.table} \\end{deluxetable}  In Table \\ref{gcbudget.table} we summarize the star clusters that we conclude may be associated with Sgr to high/moderate/low confidence, and the tidal arm in which they appear to reside. Out of 15 candidate Sgr globular clusters selected via the criterion $P_3 < 0.15$, we find that five (Arp 2, M 54, NGC 5634, Terzan 8, and Whiting 1) are very likely to be associated with the Sgr dwarf, an additional 4 (Berkeley 29, NGC 5053, Pal 12, Terzan 7) are moderately likely to be associated, two (NGC 4147, Pal 2) may be associated with relatively low confidence, and four (M 53, NGC 288, Pal 5, Pal 15) are very unlikely to be associated. Taking as our final Sgr globular cluster sample those in the high- and moderate-confidence categories, we conclude that 5-9 globular clusters are genuinely associated with Sgr while 6-10 simply appear to be associated by chance. This conclusion is consistent with our expectation, based on statistical realizations of a randomly distributed artificial globular cluster population, that $8 \\pm 2$ of the 15 candidates would be genuine Sgr clusters. As illustrated by Table \\ref{prevcands.table}, the clusters which we conclude are most likely to be associated with Sgr are generally those for which there was the most consistent agreement on Sgr membership in the literature.  We are able, however, to rule out various candidates such as M 2, NGC 288, and NGC 6356 which have also been proposed by various authors.  Our overall {\\it number} of Sgr-member globular clusters is consistent with Bellazzini et al. (2003a), who concluded (at 95\\% confidence) that at least eight globular clusters are likely to be associated with the Sgr system; four `outer-halo' globular clusters in addition to the four `core' clusters M 54, Arp 2, Terzan 7, and Terzan 8.   Curiously,  of the 5-9 globular clusters which we associate with the Sgr stream only one or two (Berkeley 29, and possibly Pal 12) lie in the leading arm of the Sgr stream. In part, this may be due to selection bias; $\\sim75$\\% of the clusters assigned to the T1 arm lie within $\\sim100^{\\circ}$ of the Sgr dwarf.  The corresponding section of the L1 arm lies almost entirely in the zone of avoidance at low Galactic latitudes where sparse clusters such as Whiting 1 and Arp 2 would be more difficult to detect. However, this may also point to a difference in the survivability of clusters in the leading versus trailing Sgr tidal arms. Since clusters in the leading arm would pass more closely to the Galactic center than clusters in the trailing arm, gravitational shocks from the Galactic disk may be more efficient in disrupting clusters in the leading arm.  Indeed, as illustrated in Figure \\ref{survival.fig}, sparse clusters such as Whiting 1 in particular lie significantly outside the survival regime for orbits 12 kpc from the Galactic center (roughly the perigalacticon of the Sgr leading arm tidal debris on some orbits in the triaxial halo model of LM10), and may therefore be disrupted relatively quickly once they are removed from Sgr into the leading tidal stream.    \\begin{figure} \\plotone{f11.eps} \\caption{Vital diagram (after Gnedin \\& Ostriker 1997) for high, moderate, and low-confidence Sgr globular clusters.  Masses and half-light radii for the classical globular clusters are taken from tabulations by Gnedin \\& Ostriker (1997) and Harris (1996).  These parameters are estimated for Berkeley 29 and Whiting 1 using data from Kaluzny (1994), Tosi et al. (2004), and Carraro (2005).  Overlaid in black solid lines are the survival regions for distances $r_{\\rm GC} =$ 12, 8, and 5 kpc from the Galactic center in the Ostriker \\& Caldwell (1983) galaxy model.} \\label{survival.fig} \\end{figure}      It is perhaps unsurprising that  we are unable to associate any  stellar clusters  with the older T2/L2 wraps of Sgr tidal debris confidently, given the possible destruction of these clusters over time, the current lack of evidence for these sections of Sgr stream in the Galactic halo and the correspondingly poor constraints on the LM10 model in this regime. Our null result may, however, constitute further evidence that the Sgr dwarf has only been losing appreciable quantities of stellar debris in the last $\\sim$ 3 Gyr, which points to a constraint on the interaction time of the dwarf similar to that currently implied by wide-field observations of the stellar debris streams.  Alternatively, discrete stellar systems such as globular clusters may simply have formed sufficiently deep in the Sgr gravitational potential well that they have only been stripped from the dwarf after it has already experienced significant tidal disruption.    We find no compelling evidence to associate Sgr with any of the population of ultra-faint Milky Way satellites recently discovered in the SDSS. While some of the ultra-faint satellites (e.g., Boo II, Segue 1) are roughly co-spatial with sections of the Sgr stream and have similar radial velocities, this apparent correlation may be a product of the survey area in which the satellites were discovered; the ultra-faint satellite population is statistically consistent with being unassociated with the Sgr tidal streams. It is not possible for us to discriminate between the cases in which some of these systems (a) formed within Sgr, (b) were loosely associated with Sgr and fell into the Milky Way as a group (e.g.,  Li \\& Helmi 2008; D'Onghia \\& Lake 2009), or (c) are completely unrelated to Sgr. We are therefore unable to cast any light upon the origins of these enigmatic systems, except that they were  not likely previously tightly bound parts of the Sgr dSph.   Niederste-Ostholt et al. (2010) have recently provided an estimate of the absolute magnitude of the initial Sgr dwarf by summing the light contained in the remnant Sgr core and the tidal streams.  Adopting the fainter of their two luminosity functions (which is most consistent with the LM10 assumption that the present absolute magnitude of the Sgr core $M_V = -13.64$), assuming that the tidal tails are symmetric in mass content, and assuming that there is negligible luminosity contained in regions of the stream for which there is no observational evidence in the combined 2MASS + SDSS surveys, the Niederste-Ostholt et al. (2010) analysis indicates that the original Sgr dwarf had $M_V \\approx -15.0$ (with $\\sim 70$\\% of this light now residing in the tidal tails). Our 5-9 globular clusters therefore correspond to a specific globular cluster frequency $S_N = 5 - 9$ for the original Sgr dSph, where $S_N = N_{\\rm GC} \\times 10^{0.4 (M_V + 15)}$ (Harris \\& Van den Bergh 1981). While this frequency is high compared to that of the Milky Way, LMC, and SMC ($S_N = 0.7, 0.8, 1.0$ respectively based on the compilation by Forbes et al. 2000), it is substantially less than that of Fornax ($S_N = 29$) and of previous estimates of the specific frequency of Sgr ($S_N \\sim 25$; e.g., Forbes et al. 2000; Forbes \\& Bridges 2010),\\footnote{The main reason for the discrepancy between our estimate of the specific frequency of Sgr and that of previous studies is our assumption of a significantly brighter original satellite.} and similar to that of many dE galaxies (Strader et al. 2006; Brodie \\& Strader 2006, and references therein).     We caution, however, that our analysis is subject to numerous uncertainties, the greatest of which is the uncertainty in the LM10 model of the Sagittarius tidal streams. While the LM10 model is reasonably well constrained for the most recent 3 Gyr of tidal debris for which there is the strongest evidence to date in the Galactic halo, its predictions for the possible locations and properties of earlier epochs of tidal debris are only extrapolations.  If significant evolution has occurred in the orbit of Sgr over time (e.g., via dynamical friction) it may be difficult to identify debris (whether stars or larger stellar systems) from the earliest epochs whose present orbits may not resemble the current orbit of Sgr.  Our inventory of the substructures associated with Sgr may therefore be an underestimate. In future, proper motions of the quality expected to be provided by NASA's Space Interferometry Mission (SIM) are expected to permit us to assess the case for membership of individual stellar systems with the Sgr tidal stream with greater confidence.   \\subsection{The Age-Metallicity Relation and Horizontal Branch Morphology of Sgr Globular Clusters}  In Figure \\ref{amr.fig} (left panel), we plot the ages and metallicities (drawn from tabulations by Marin-Franch et al. 2009, and Forbes \\& Bridges 2010) of our eleven high/moderate/low confidence Sgr globular clusters. As discussed recently by Forbes \\& Bridges (2010) for a sample of Sgr-candidate clusters drawn from the literature, the clusters describe a relatively well defined age-metallicity relation (AMR), which we note is generally consistent with the closed-box Sgr AMR derived by Siegel et al. (2007). We identify three groupings of clusters in particular in Figure \\ref{amr.fig}: M 54\\footnote{Note that the age and metallicity of M 54 was adopted from the Siegel et al. (2007) M 54 MPP measurement.}, NGC 5053, NGC 5634, Pal 2, and Ter 8 resemble the oldest Sgr stellar population (`M54 MPP'), Arp 2 and NGC 4147 resemble the second-oldest Sgr stellar population (`Sgr MPP'), and Be 29, Pal 12, Ter 7, and Whiting 1 resemble the intermediate age Sgr stellar population (`SInt').    In Figure \\ref{amr.fig} (right panel) we show the Lee diagram of the candidate Sgr globular clusters, plotting metallicity versus horizontal branch (HB) morphology $HBR = (B-R)/(B+V+R)$, where $B$, $V$, and $R$ represent the number of blue, variable, and red horizontal branch stars respectively. As has often been noted previously (e.g., Mottini \\& Wallerstein 2008, and references therein), the Sgr candidate globular clusters are a bimodal group of red- and blue-HB clusters. Intriguingly, the strongest Sgr cluster candidates (i.e., cyan filled squares) are those which, despite their $\\sim 6$ Gyr spread in ages, show the least evidence for a second-parameter affecting their HB morphologies and lie closest to the single-parameter (i.e., metallicity) locus defined by Zinn (1993a; solid line in the right panel of Figure \\ref{amr.fig}).\\footnote{While M 54 lies to the left of the single-parameter locus, this is likely due to confusion between the blue HB  of M 54 and the red HB of the more metal-rich Sgr stellar populations.} The moderate-confidence clusters Be 29, Pal 12, and Ter 7 also follow the single-parameter locus, while NGC 5053 and the two low-confidence clusters NGC 4147 and Pal 2 have HBs redder than expected for their metallicity. The absence of a second-parameter effect in the Sgr globular clusters is in marked contrast to the globular cluster systems of the LMC and Fornax (e.g., Zinn 1993b; Buonanno et al. 1999; Mackey \\& Gilmore 2004), which show moderate/strong influences of a second parameter on their respective HB types at a given metallicity. This may suggest that the strength of the second-parameter effect on the HB  morphologies of accreted globular clusters could somehow be determined by the subhalo in which they were formed, offering a criterion by which to associate clusters in the halo of the Milky Way with their parent satellites.      Although there is a strong gradient in the chemistry of stars along the Sgr streams (see, e.g., Chou et al. 2007, 2010), there is no comparable gradient for the Sgr globular clusters. While the two globular clusters wrapped farthest from the Sgr dwarf (NGC 5053 and NGC 5634, wrapped $266^{\\circ}$ and $293^{\\circ}$ from Sgr respectively) are the most metal-poor of the clusters, their closest chemical kin (Arp 2, M 54, Ter 8) lie $< 10^{\\circ}$ from Sgr. Similarly, the  relatively metal-rich clusters Whiting 1  and Berkeley 29 lie in the dynamically young  ($103^{\\circ}$ separation from Sgr) and old ($183^{\\circ}$ separation from Sgr) sections of the stream respectively. Indeed, if Berkeley 29 formed 4.5 Gyr ago (Carraro et al. 2004; 3.4 - 3.7 Gyr ago from the estimate of Tosi et al. 2004) in the `SInt' burst of star formation, the LM10 model suggests that it must have been stripped from Sgr less than a Gyr after its formation since it would have taken $\\sim 3-4$ Gyr since stripping to reach its present longitudinal separation from Sgr.   \\begin{figure*} \\plotone{f12.eps} \\caption{Left panel: Age-metallicity relation for candidate Sgr globular clusters (where observational data is available). Filled cyan squares represent high confidence candidates, green crosses moderate confidence candidates, and red open boxes low confidence candidates. Overplotted is the AMR of Siegel et al. (2007; solid line), which best fits their reconstructed star formation history of Sgr (elliptical shaded regions). Right panel: Horizontal branch morphology for candidate Sgr globular clusters as a function of metallicity.  Overlaid (solid line) is the `first-parameter' locus defined by Zinn (1993a), objects to the left of this line are those with a significant second-parameter effect shaping their horizontal branch morphologies.} \\label{amr.fig} \\end{figure*}"
    },
    "1005/1005.0260_arXiv.txt": {
        "abstract": "CCD photometric observations in $VRI$ colors and spectroscopic observations of the newly discovered eclipsing binary GSC~2314-0530 (NSVS~6550671) with dMe components and very short period of $P=0.192636$ days are presented. The simultaneous light-curve solution and radial velocity solution allows to determine the global parameters of GSC~2314-0530: $T_{1}=3735$ K; $T_{2}=3106$ K; $M_{1}=0.51$ M$_{\\sun}$; $M_{2}=0.26$ M$_{\\sun}$; $R_{1}=0.55$ R$_{\\sun}$; $R_{2}=0.29$ R$_{\\sun}$; $L_{1}=0.053$ L$_{\\sun}$; $L_{2}=0.007$ L$_{\\sun}$; $i=72.5\\degr$; $a=1.28$ R$_{\\sun}$; $d=59$ pc. The chromospheric activity of its components is revealed by strong emission in the H$\\alpha$ line (with mean $EW=5\\ {\\rm \\AA}$) and observed several flares. Empirical relations mass--$M_{\\rm {bol}}$, mass--radius and mass--temperature are derived on the basis of the parameters of known binaries with low-mass dM components. ",
        "introduction": "\\label{sec:intro}  Although the M dwarfs are the most numerous stars in our Galaxy, the mass, metalicity and age dependencies of their stellar luminosities and radii are poorly calibrated. The reason is the selection effect that plays against the detection of fainter and smaller stars.  Less than 20 binaries with low-mass dM components have empirically-determined masses, radii, luminosities and temperatures (see Section \\ref{sec:global}, Table \\ref{tab:stars}). As a result the mass-luminosity relation is determined by only a few low-mass stars. This deficiency hindered the development of the models for the cool dense atmospheres of the M dwarfs. It is established that all available models underestimate the radii (by around 10--15 per cent) and overestimate the temperatures (by 200--300 K) of short-period binaries with dM components \\citep{ribas03,maceroni04}.  The Northern Sky Variability Survey (NSVS) contains a great number of photometric data \\citep{wozniak04} that allows searching of variable stars and determination of their periods and types of variability. A multiparametric method for search for variable objects in large datasets was tested on the NSVS \\citep{dimitrov09} and as a result many eclipsing stars were discovered. One of them was GSC~2314-0530 $\\equiv$ NSVS~6550671 ($\\alpha$=02$^{\\rm h}20^{\\rm m}50\\fs9$, $\\delta$=+$33\\degr 20\\arcmin 46\\farcs6$).  On the base of the NSVS photometry obtained in 1999--2000 we derived the ephemeris:  \\begin{equation}\\label{equ:nsvs} HJD({\\rmn {MinI}})=2451352.062 + 0.192637 \\times E \\end{equation} and built its light curve (Fig. \\ref{fig:nsvs}).  \\begin{figure} \\centering \\includegraphics[width=0.75\\columnwidth]{fig1.eps} \\caption{NSVS photometry of GSC~2314-0530} \\label{fig:nsvs} \\end{figure}  We found that this star has been assigned also as SWASP~J022050.85+332047.6 according to the SuperWASP photometric survey \\citep{pollacco06}. \\citet{norton07} reported its coincidence with the {\\it ROSAT} X-ray source 1RXS~J022050.7+332049.  Initially GSC~2314-0530 attracted our interest by its short orbital period because there were only several systems with non-degenerate components and periods below the short-period limit of 0.22 days \\citep{rucinski07}: GSC~1387-0475 with $P=0.217811$ d \\citep{rucinski07,rucinski08}, ASAS~J071829-0336.7 with $P=0.211249$ d \\citep{pribulla09}, the star V34 in the globular cluster 47 Tuc with $P=0.2155$ d \\citep{weldrake04} and BW3~V38 with orbital period $P=0.1984$ d \\citep{maceroni97,maceroni04}.  When we established that the components of GSC~2314-0530 were dM stars our interest increased and we undertook intensive photometric and spectral observations in order to determine its global parameters and to add a new information for the dM stars as well as for the short-period binaries.  \\begin{figure} \\centering \\includegraphics[width=0.75\\columnwidth]{fig2.eps} \\caption{Observed field around GSC~2314-0530} \\label{fig:chart} \\end{figure} ",
        "conclusions": "  (1) This star is the shortest-period binary with dM components which period is below the short-period limit.  (2) By simultaneous radial velocity solution and light curve solution we determined the global parameters of GSC~2314-0530: inclination $i=72.5\\degr$; orbital separation $a=1.28$ R$_{\\sun}$; masses $M_{1}=0.51$ M$_{\\sun}$ and $M_{2}=0.26$ M$_{\\sun}$; radii $R_{1}=0.55$ R$_{\\sun}$ and $R_{2}=0.29$ R$_{\\sun}$; temperatures $T_{1}=3735$ K and $T_{2}=3106$ K; luminosities $L_{1}=0.053$ L$_{\\sun}$ and $L_{2}=0.007$ L$_{\\sun}$; distance $d=59$ pc.  (3) We derived empirical relations mass--$M_{\\rm {bol}}$, mass--radius and mass--temperature on the basis of the parameters of known binaries with low-mass dM components.  (4) The distorted light curve of GSC~2314-0530 were reproduced by two cool spots on the primary component. The next sign of the activity of GSC~2314-0530 is the strong H$\\alpha$ emission of its components. Moreover we registered 6 flares of GSC~2314-0530. Half of them occurred at the phases of maximum visibility of the larger stable cool spot on the primary.  The analysis of all appearances of magnetic activity revealed existence of long-lived active area on the primary of GSC~2314-0530. The high activity of the target is natural consequence of the fast rotation and low temperatures of its components.  Our study of the newly discovered short-period eclipsing binary GSC~2314-0530 presents a next small step toward understanding dMe stars and adds a new information to the poor statistic of the low-mass dM stars. Recently they became especially interesting as appropriate targets for planet searches due to the relative larger transit depths."
    },
    "1005/1005.0395_arXiv.txt": {
        "abstract": "Continuous gravitational-wave (CW) signals such as emitted by spinning neutron stars are an important target class for current detectors. However, the enormous computational demand prohibits fully coherent broadband all-sky searches for prior unknown CW sources over wide ranges of parameter space and for yearlong observation times. More efficient hierarchical ``semicoherent\" search strategies divide the data into segments much shorter than one year, which are analyzed coherently; then detection statistics from different segments are combined incoherently. To optimally perform the incoherent combination, understanding of the underlying parameter-space structure is requisite. This problem is addressed here by using new coordinates on the parameter space, which yield the first analytical parameter-space metric for the incoherent combination step. This semicoherent metric applies to broadband all-sky surveys (also embedding directed searches at fixed sky position) for isolated CW sources. Furthermore, the additional metric resolution attained through the combination of segments is studied. From the search parameters (sky position, frequency, and frequency derivatives), solely the metric resolution in the frequency derivatives is found to significantly increase with the number of segments. ",
        "introduction": "}  Direct detection of gravitational waves would not only validate Einstein's theory of general relativity but also constitute an important new astronomical tool. Continuous gravitational-wave (CW) signals are expected, for instance, from rapidly rotating neutron stars through various emission mechanisms \\cite{owen-1998-58,ushomirsky:2000,cutler:2002-66,jones-2002-331,owen-2005-95}. Most such stars are anticipated to be electromagnetically invisible, but might be detected and studied via gravitational waves. With current Earth-based detectors, such as LIGO~\\cite{ligoS5}, numerous efforts are underway to search for CW sources \\cite{pshS4:2008,S4EAH,S5R1EAH,eahurl}, and observational upper limits have already allowed one to constrain the physics of neutron stars~\\cite{crab-2008,powerflux:2009}.  The expected CW signals are extremely weak. Hence detection requires very sensitive data analysis techniques to extract these signals from detector noise. In the work of~\\cite{jks1} a powerful method has been derived which is based on the principle of maximum likelihood detection, leading to coherent matched filtering. The CW signals are quasimonochromatic with a slowly changing intrinsic frequency. For a terrestrial detector, the Earth's motion relative to the solar system barycenter (SSB) generates a Doppler modulation in amplitude and phase of the waveform.  As shown in~\\cite{jks1}, the parameters describing the signal's amplitude variation may be analytically eliminated by maximizing the coherent matched-filtering statistic.  Thus, one only searches over the remaining parameters describing the signal's phase: the source's sky location, frequency, and frequency derivatives (``spindowns''). The so-obtained coherent detection statistic is called the \\mbox{$\\F$-statistic}, which can also be generalized to employ multiple detector data streams \\cite{CutlerMultiIFO}.  Finite computational resources are what imposes severe limits on the sensitivity of broadband all-sky searches for prior unknown CW sources~\\cite{jks1,bccs1:1998}. In the fully coherent \\mbox{$\\F$-statistic} approach, one must convolve the full data set with signal waveforms (templates) corresponding to all possible sources. But the number of templates required to cover the search parameter space increases as a high power of the coherent integration time. For year-long data sets, searching a realistic range of parameter space would demand more computing power than available on Earth~\\cite{bccs1:1998,jks1}.  Therefore, a fully coherent search is restricted to much shorter integration times.  A more efficient analysis of data sets which span yearlong time intervals is achieved via less expensive hierarchical semicoherent methods \\cite{bc2:2000,cutler:2005,hough:2005,pletsch:GCT}. In such a method, the data are divided into segments of duration~$T$, where~$T$ is much smaller than one year. This allows one to use a \\emph{coarse} grid of templates, on which the \\mbox{$\\F$-statistic} is calculated for each segment. Then the $\\F$-statistics from all segments (or statistics derived from~$\\F$) are incoherently combined using a common \\emph{fine} grid of templates. Since phase information is discarded between segments, this latter step is called incoherent and thus the search methodology as a whole is referred to as \\emph{semicoherent}.  An important long-standing problem in these semicoherent methods has been the design of, and link between, the coarse and fine grids. To address this problem it is essential to understand the underlying parameter-space structure. In this context, the geometric approach has been proven to be very useful, introducing a \\emph{metric} on parameter space, as  first done in \\cite{Sathy1:1996,owen:1996me}. The key quantity in this respect is called \\emph{mismatch}~$\\mm$, which is the fractional loss in expected $\\F$-statistic (or sum of expected \\mbox{$\\F$-statistics} in the incoherent step) for a given signal $\\vDoppler_\\sig$ at a nearby grid point~$\\vDoppler$ in phase-parameter space~$\\Pspace$.  Taylor-expanding $\\mm$ (to quadratic order) at $\\vDoppler_\\sig$ in the differences of $\\vDoppler_\\sig$ and $\\vDoppler$ defines the parameter-space metric.  While the parameter-space metric for coherent searches has been examined in detail \\cite{bccs1:1998,ptolemetric,prix:2007mu}, the metric for semicoherent searches is comparably much less well studied. However, the literature does exhibit several quests for a semicoherent metric \\cite{cutler:2005,hough:2005,Dhurandhar2008,Watts2008}. About a decade ago, the first general discussion of the semicoherent metric for CW searches was given in~\\cite{bc2:2000}, along with numerical investigations. But only recently, the first fully analytical semicoherent metric has been found, leading to a significantly improved CW search method~\\cite{pletsch:GCT}. This recent progress is based on a better understanding of the global parameter-space correlations~\\cite{pletsch:2008}, which were first examined in~\\cite{prixitoh:2005}. In turn, this insight provides new coordinates on parameter space, enabling the analytical calculation and study of the semicoherent metric.  The present paper extends the recent work of~\\cite{pletsch:GCT} to greater generality and provides the essential technical basis for a parameter-space metric formalism for semicoherent CW searches. Additionally, complete and fully analytic semicoherent metric results are presented, which are ready to use for practical implementations of semicoherent searches, serving the earlier-mentioned quests of previous literature. The results apply to broadband all-sky surveys (embedding directed searches with fixed sky position) for isolated CW sources.  Section~\\ref{sec:MatchedFilter} briefly reviews matched filtering for CW signals and the parameter-space metric formalism in general. Section~\\ref{sec:SemiCohSearches} elucidates how the metric is obtained for semicoherent searches. To evaluate and study the semicoherent metric, new coordinates on parameter space are defined in \\Sref{sec:NewCoords}. In \\Sref{sec:OneSpindown}, the semicoherent metric is derived and investigated for the case of CW signals whose intrinsic frequency changes linearly with time (parametrized by one spindown parameter). In particular, the refinement factor is introduced as the ratio of the number fine-grid and coarse-grid templates, quantifying the additional parameter-space metric resolution due to combination of many segments. Section~\\ref{sec:TwoSpindowns} extends these results to the case of CW signals whose intrinsic frequency can change quadratically with time (considering two spindown parameters). Finally, a short conclusion follows in \\Sref{sec:Conclusion}. ",
        "conclusions": "}  A formalism has been presented for the parameter-space metric of semicoherent CW searches, where the data are divided into segments that are coherently analyzed and subsequently combined incoherently. By using new coordinates on parameter space, the first fully analytical semicoherent metric for broadband all-sky CW surveys has been derived. Additionally, in the new coordinates the components of both the coherent and the semicoherent metric tensor are constant, being explicitly independent of the coordinates. This entails great convenience regarding practical aspects of semicoherent CW searches.  Explicit analytic expressions of the semicoherent metric tensor components have been obtained for two typical search parameter spaces of current practical interest. First, the one-spindown case has been considered, restricting to linear frequency drifts with time as done in many current all-sky searches. Second, the semicoherent metric also has been calculated and examined for the two-spindown case, where additionally quadratic changes in frequency are taken into account.  Analytic signal template counting formulas have been provided for the coherent stage (coarse grid) as well as for the incoherent combination step (fine grid). In this respect, a useful quantity, called the \\emph{refinement factor}, has been introduced as the ratio of the number of fine-grid and coarse-grid templates. Thus, the refinement factor describes (coordinate independently) the additional parameter-space metric resolution gained from the combination of segments. Moreover, the scaling of the refinement factor with the number of segments has been found to be predominantly determined solely by the spindown parameters.  The present results also embed the case of directed semicoherent searches, where the sky position is known and hence is not a search parameter. Thereby, the search parameter space consists only of frequency and spindown parameters. The corresponding semicoherent metric tensor components are identical to ones derived here. The resulting scalings of the refinement factor with the number of segments also hold, since governed by the spindown parameters as described above.  The formalism presented in this paper assumes segments that are very short compared to one year. The Earth's orbital motion then can be described by a low-order Taylor expansion, and therefore be modeled by changes in frequency and frequency derivatives. In the future, however, increased computing power might allow one to use segments  substantially longer than a few days and hence issues remain to be explored in such circumstances.   \\newpage"
    },
    "1005/1005.5265_arXiv.txt": {
        "abstract": "Warm cores (or hot corinos) around low-mass protostellar objects show a rich chemistry with strong spatial variations. This chemistry is generally attributed to the sublimation of icy mantles on dust grains initiated by the warming effect of the stellar radiation. We have used a model of the chemistry in warm cores in which the sublimation process is based on extensive laboratory data; these data indicate that sublimation from mixed ices occurs in several well-defined temperature bands. We have determined the position of these bands for the slow warming by a solar-mass star. The resulting chemistry is dominated by the sublimation process and by subsequent gas-phase reactions; strong spatial and temporal variations in certain molecular species are found to occur, and our results are, in general, consistent with observational results for the well-studied source IRAS 16293-2422. The model used is similar to one that describes the chemistry of hot cores. We infer that the chemistry of both hot cores and warm cores may be described by the same model (suitably adjusted for different physical parameters). ",
        "introduction": "The early stages of low-mass star formation are observed at millimetre and submillimetre wavelengths, where both the dust continuum as well as molecular emission from the outer regions of the envelope can be detected.  The dynamical aspect of low-mass star formation, with low mass stars forming from dense and cold condensations called prestellar cores, and evolving into Class 0 and I sources, is relatively well defined. What is less well understood is how prestellar cores evolve chemically into protostellar cores which then, eventually, develop protoplanetary disks.  Nevertheless, it is now clear that, as in the case of massive star formation, towards the end of the collapse phase the gas and dust are dense and cold enough that (other than H$_2$ and He) material has frozen onto the dust and that surface chemistry has taken place, leading to a rich gas phase chemistry once the star heats the surrounding dust so that icy mantles returned to the gas phase. The warm cores where the icy mantles sublime are often called `hot corinos' (e.g. \\citealt{cec04asp,cec07msl}) due to their chemical resemblance with hot cores. Recent observations toward several warm cores around low-mass solar-like protostars have shown that they are rich in large, complex organic molecules (COMs) including methanol CH$_3$OH, methyl formate HCOOCH$_3$, methyl cyanide CH$_3$CN, and more complex species such as ethylene glycol (CH$_2$OH)$_2$ (e.g. \\citealt{caz03,bot04,hol04,bis08}). For a comprehensive review on complex organic interstellar molecules we refer the reader to \\citet*{her09}. % In addition, they show some interesting deuterated species (e.g. \\citealt{cec98a,caux05}), and long-chain unsaturated hydrocarbons and cyanopolyynes \\citep{sak08}. In fact, the envelope of Class 0 sources can be divided chemically into two zones: the zone where all the ices sublime (dust temperature $\\ge 100$ K), i.e. the warm core region, and the outer cold envelope (dust temperature $\\le 100$ K) (e.g. \\citealt{cec05iau}). The chemistry in the outer envelope is similar to that in pre-stellar cores, which are cold ($\\sim 10$ K) and dense ( $\\le 1\\times10^{5}$ cm$^{-3}$) condensations believed to be precursors of the Class 0 sources \\citep{ber07}. Many molecules in the outer envelope are frozen onto the grain mantles, so that their abundances are often lower than those in molecular clouds. In the inner regions, sublimation is clearly important.  The first model of the thermal structure of a low-mass protostellar envelope was developed by \\citet{cec96}. In this model, the authors considered the chemistry and resulting line emission within the inside-out collapse model \\citep{shu77} but down to much smaller scales of $\\sim 10$ AU. In their model, the gas temperature in the innermost regions of the envelope reaches $\\sim 100$ K resulting in the sublimation of H$_2$O from grain mantles.  More recent models (e.g. \\citealt{cec20b, mar04}) suggest that molecular emission (e.g. from H$_2$CO) is emitted by two components: a cold H$_2$CO-poor outer envelope and a warm H$_2$CO-rich core. The model calculations show a dramatic increase in the abundances of H$_2$CO, H$_2$O, and SiO in the inner, warmer, and denser region of the envelope. This abrupt change occurs when the dust temperature exceeds 100 K and grain mantles sublime.  Evidently, in situations where a cool outer envelope and a warmer core exist, sublimation - and the subsequent gas-phase chemistry that this causes - will be important. Also, the sublimation will be a time- and space-dependent event as the central protostar warms up. Given the role of sublimation, it is important to take account of recent laboratory studies of the sublimation of ices \\citep{col03a,col04} which show that different species enter the gas phase at different temperatures, and that the sublimation occurs not continuously through a gradual warm-up but in several well-defined and narrow temperature bands. Such a process may give rise to a rather different gas-phase chemistry than a process in which instantaneous sublimation of all species in a mantle occurs, as found to be the case for hot cores \\citep{viti04}.  Our purpose in this paper is, therefore, to model the sublimation process, taking account of the laboratory data, in the context of warm cores in low-mass star-forming regions. We wish to assess whether the more complex sublimation mechanisms indicated by the laboratory studies significantly affect the detectable chemistry in these objects, as compared to models in which the desorption is instantaneous for all species. The model we use is developed from a similar model used for the study of hot cores in regions of massive star formation \\citep{viti04}. Further, if the two similar models can successfully describe cores in both types of star-forming region, then this may be taken as support for the view that star formation is carried out by essentially the same process, regardless of mass. We describe our model in Section 2, the results are given in Section 3, and we make some conclusions in Section 4. ",
        "conclusions": "We report here results from a chemical model of a warm core (or hot corino) around a low-mass star. The core contains a protostellar object of solar mass whose radiation warms its inner part. This leads to sublimation of icy mantles deposited during the preceding collapse phase. We have introduced a detailed model of sublimation based on extensive laboratory results of \\citet{col04}. These experiments imply that sublimation from mixed ices during the warming process occurs in several well-defined temperature bands. We have computed the positions of these bands for the solar-mass model and follow the resulting chemistry in a time- and space-dependent manner.  Although we have not attempted to model any specific astronomical source, the physical parameters we have used are sufficiently similar to those inferred from observations of a well-studied object, IRAS 16293-2422,  for us to compare the model results with observations. The predicted and inferred molecular abundances agree reasonably well.  The model that we have used is essentially a hot core model in which several physical parameters have been adjusted to be appropriate to warming by a solar-mass protostar. The generally satisfactory behaviour of the model results suggests that the chemical processes for both warm cores in low-mass star forming regions and hot cores in high mass star forming regions are similar.  Observations of warm cores show the existence at particular values of A$_V$ of very substantial `jumps' in abundance (by several orders of magnitude) in certain molecular abundances. We find that these jumps arise as a consequence of sublimation and subsequent rapid gas-phase chemistry. Such jumps do not occur in models that describe sublimation as an instantaneous process, rather than the staged process we have introduced.  We predict that ratios of certain sulphur species should be excellent tracers of physical conditions in warm cores, regardless of the (unknown) total sulphur abundance."
    },
    "1005/1005.0506_arXiv.txt": {
        "abstract": "{Secular evolution gradually shapes galaxies by internal processes, in contrast to early cosmological evolution which is more rapid. An important driver of secular evolution is the flow of gas from the disk into the central regions, often under the influence of a bar. In this paper, we review several new observational results on bars and nuclear rings in galaxies. They show that these components are intimately linked to each other, and to the properties of their host galaxy. We briefly discuss how upcoming observations, e.g., imaging from the {\\it Spitzer} Survey of Stellar Structure in Galaxies (S$^4$G), will lead to significant further advances in this area of research.} ",
        "introduction": "\\label{intro}  The general topic of galaxy evolution is enjoying widespread attention in the literature. The aim is to answer, through observations and numerical modelling, how galaxies have evolved from the earliest stages of the Universe, or from a subsequent epoch of formation, to the shape in which we observe them in the local Universe. The two main strands in this wide topic are cosmological evolution of galaxies, which deals with their formation and early evolution, and secular evolution, with which we mean the internal evolution of galaxies, under the influence of the dynamical actions of, e.g., bars or spiral arms. Secular evolution, as comprehensively reviewed by Kormendy \\& Kennicutt (2004),  is a relatively slow process compared to the more rapid evolution undergone by galaxies in the early Universe, the latter often due to mergers and galaxy-galaxy interactions.  Because the early stages of galaxy formation and evolution are hard or sometimes impossible to observe due to the combined effects of distance, redshift, and dust extinction, the detailed study of nearby galaxies is one of the very few ways to confirm the detailed predictions of models of large-scale galaxy formation and evolution. This kind of study can be nicknamed `galactic palaeontology', because in local galaxies we study the `fossil record' of billions of years of galaxy evolution---both cosmological and secular. To read and interpret this fossil record, a combination of many different observational and interpretational techniques must be used, from observations of individual stars in our own Milky Way and the most nearby galaxies to photometric and kinematic observations across many different wavelengths in external galaxies, all combined with a wide range of interpretational, analytical and numerical tools.  One of the major drivers of the internal evolution is the flow of gaseous material, from the disk to the central regions of the galaxy. To move inwards, rotating gas must lose angular momentum, and it can do so by shocking and under the influence of a non-axisymmetric potential. That in turn gets set up by a bar, by interactions or minor mergers, or even by less obvious deviations from axisymmetry like spiral arms, ovals, or lenses (e.g., Schwartz 1984; Shlosman et al. 1989, 1990; Knapen et al. 1995; Kormendy \\& Kennicutt 2004; Comer\\'on et al. 2010). Bars are very common in galaxies, about two thirds of local galaxies have a bar (de Vaucouleurs et al. 1991; Sellwood \\& Wilkinson 1993; Moles et al. 1995; Ho et al. 1997; Mulchaey \\& Regan 1997; Hunt \\& Malkan 1999; Knapen et al. 2000; Eskridge et al. 2000; Laine et al. 2002; Laurikainen et al. 2004a; Men\\'endez-Delmestre et al. 2007; Marinova \\& Jogee 2007; Sheth et al. 2008; Laurikainen et al. 2009). As we will see below, in Sect.~\\ref{rings}, nuclear rings can occur in unbarred galaxies, and they seem to prove that the non-axisymmetry induced by a spiral or oval may well be enough to induce gas inflow, and thus lead to secular evolution. Although most galaxies are barred, secular evolution is thus not dependent on the presence of a bar.  Galactic interactions are rare, occurring in only about 2\\% of local galaxies (up to 4\\% if merely bright galaxies are considered; Knapen \\& James 2009). Close companions to local galaxies are much more common, and Knapen \\& James (2009) found that some 15\\% of local galaxies have a companion not more than 3\\,mag fainter than itself within a radius of five times the diameter of the galaxy under consideration, and within a range of $\\pm200$\\,km\\,s$^{-1}$ in systemic velocity. The effects on the star formation rate in galaxies of the presence of a close companion, and even of interactions, is, perhaps surprisingly, limited (observations by Bushouse 1987; Smith et al. 2007; Woods \\& Geller 2007; Li et al. 2008; Knapen \\& James 2009; Jogee et al. 2009; Rogers et al. 2009; Ellison et al. 2010, corroborated by numerical simulations by Mihos \\& Hernquist 1996; Kapferer et al. 2005; Di Matteo et al. 2007, 2008; Cox et al. 2008). Statistically, the star formation rate is raised by a factor of just under two by the presence of a close companion, but the H$\\alpha$ equivalent width is hardly increased at all (Knapen \\& James 2009). This implies that even though galaxies with close companions tend to form stars at a higher rate, they do so over extended periods of time, and not as a burst. Even the majority of the truly interacting galaxies in Knapen \\& James's (2009) sample have unremarkable star formation properties. The reason that extreme star-forming galaxies, such as ultra-luminous infrared galaxies (ULIRGs) are found to be almost ubiquitously interacting must surely be a selection effect. In general, interactions do not always cause starbursts, and starbursts do not always occur in interacting galaxies.  In fact, Knapen \\& James (2009) used their sample of 327 nearby disk galaxies to explore how one might best define the term `starburst'. They concluded that {\\it none} of the definitions that are in common use in the literature can be considered to be objective and generally discriminant. For instance, selecting galaxies on the basis of their high star formation rate yields large star-forming disk galaxies. Selecting those with high equivalent widths (apparently a {\\it bona fide} starburst discriminator as this selects galaxies with a much enhanced current star formation rate as compared to the average rate in the past), yields primarily late-type galaxies of very small mass, whose star formation activity is caused by one or a few H{\\sc ii} regions (and which will have very low impact on the intergalactic medium through, e.g., stellar winds). And selecting galaxies with the shortest gas depletion timescales does not only select galaxies with very high current star formation rates, but also gas-poor early-type galaxies with a very small star formation rate. The conclusion of Knapen \\& James (2009) is that starbursts are very hard to define properly, and the use of the term should be restricted to well-described small numbers of objects.  This review deals with aspects of secular evolution, and how we can trace its actions back through the detailed study of structural components, particularly bars and rings, in nearby galaxies. The interpret the effects of these agents, the tools we will use here are primarily optical and infrared imaging and two-dimensional kinematic mapping. Other authors have presented reviews on galactic evolution, and in particular the paper by Kormendy \\& Kennicutt (2004) presents an authoritative review of the subject of secular evolution. We will supplement that by presenting selected recent results on bars and rings that highlight the intricate and detailed connections that exist between the different structural components of a galaxy and the overall galactic evolution.  In Section~\\ref{bars} of this paper, we will review how the strength of a bar is connected to many of the basic properties of the bar, such as its length, or the shape of its dust lanes. We will also see that bars are indeed connected to spirals, and how bars in S0 galaxies may be different from those in spirals. Section~\\ref{rings} describes how the basic physical properties of the host galaxy and, where present, its bar condition the location and morphology of a nuclear ring, thus highlighting the close physical connections between these components. Section~\\ref{future} discusses how the {\\it Spitzer} Survey of Stellar Structure in Galaxies (S$^4$G; Sheth et al. 2010) will deliver the data which should allow us to make significant further progress in the study of galaxy evolution by means of detailed analyses of the stellar component in a large sample of nearby galaxies. We briefly present our conclusions in Section~\\ref{conclusions}. ",
        "conclusions": ""
    },
    "1005/1005.5323_arXiv.txt": {
        "abstract": "{ Despite being very successful in explaining the wide range of precision experimental results obtained so far, the Standard Model (SM) of elementary particles fails to address two of the greatest observations of the recent decades: tiny but nonzero neutrino masses and the well-known problem of missing mass in the Universe. Typically the new models beyond the SM explain only one of  these observations. Instead, in the present article, we take the view that they both point towards the same new extension of the Standard Model. The new particles introduced are responsible simultaneously for neutrino masses and for the dark matter of the Universe. The stability of dark matter and the smallness of neutrino masses are guaranteed by a $\\U{1}$ global symmetry, broken to a remnant $\\mathbb{Z}_2$. The canonical seesaw mechanism is forbidden and neutrino masses emerge at the loop level being further suppressed by the small explicit breaking of the $\\U{1}$ symmetry. The new particles and interactions are invoked at the electroweak scale and lead to  rich phenomenology in colliders, in lepton flavour violating rare decays and in direct and indirect dark matter searches, making the model testable in the coming future.}  \\begin{document} ",
        "introduction": "Despite all its triumphs, the Standard Model  (SM) of elementary particles fails to explain two of the greatest observations of the recent decades: tiny but nonzero neutrino masses and the missing mass of the Universe commonly explained by Dark Matter (DM). Various models beyond the Standard Model (BSM) have been developed to explain each of these mysteries separately. On the contrary these two phenomena may be linked and explained within a single scenario~\\cite{Krauss:2002px,*Cheung:2004xm,*Asaka:2005an,*Ma:2006km,*Kubo:2006yx,*Chun:2006ss,*Hambye:2006zn,*Kubo:2006rm,*Aoki:2008av,*Aoki:2009vf,*Bajc:2010qj,Boehm:2006mi,Farzan:2009ji}. In the present article we will explore this possibility and propose an extension of the Standard Model in which the new particles and interactions are simultaneously responsible for the dark matter of the Universe and neutrino masses. We call our model AMEND which stands for ``A Model Explaining Neutrino masses and Dark matter''. We introduce a $\\U{1}_X$ global symmetry which is broken to a remnant $\\mathbb{Z}_2$. The latter symmetry distinguishes the new particles from the SM ones and is responsible for the stability of dark matter. At the same time, it forbids the canonical seesaw mechanism and left-handed neutrinos do not acquire a Dirac mass. Neutrino masses arise at the loop level and are suppressed by small terms which explicitly break the $\\U{1}_X$ symmetry.   More concretely, we introduce   two additional  electroweak doublets $R$ and $R^\\prime$ with opposite hypercharges. These Weyl fermions together can be  regarded as a Dirac four component spinor. Moreover, we add an electroweak scalar triplet $\\Delta$ and a complex singlet $\\phi$ which do not acquire vacuum expectation values (VEVs). The lightest neutral scalar, which is mainly the singlet $\\phi$ with a small admixture of the neutral component of $\\Delta$, will play the role of the DM particle. We expect the mass value for DM to be around the electroweak scale so the DM in our model can be categorised as a weakly interacting massive particle (WIMP). Smaller values can be obtained if fine-tuning is allowed in the DM mass term. The Yukawa coupling involving $\\Delta$ is the lepton number violating interaction which leads to Majorana neutrino masses. The neutrino masses are generated at one loop and are proportional to the mass of the new fermionic doublet which needs to be at the electroweak scale or higher. The model presented in this paper belongs to the class of models for which the neutrino mass is generated radiatively (see {\\it e.g.} \\cite{Zee:1980ai,*Babu:1988ki,*Ma:1998dn}). The low neutrino mass scale therefore requires a further suppression, in addition to the loop-factor, by either suppressing the Yukawa coupling  or by requiring a cancellation  between the various contributions. In our model both of these requirements are simultaneously enforced by a continuous Abelian symmetry $\\U{1}_X$ which forbids the  neutrino masses. Its explicit breaking to a residual $\\mathbb{Z}_2$ symmetry leads to small neutrino masses and a stable DM particle at the same time.  In our model, DM is produced thermally in the Early Universe. We  explore the various allowed annihilation channels which can proceed via $R$--$R^\\prime$, $Z$ or Higgs exchange, in order to reproduce the observed amount of dark matter $\\Omega_{\\mathrm{DM}} h^2 = 0.1131\\pm0.0034$~\\cite{Komatsu:2008hk}. We find that in the parameter range of interest, the dark matter annihilation through $s$-channel Higgs exchange dominates over the other channels. Dark matter is being searched for both directly, looking for the recoil of nucleons/electrons due to DM scattering in detectors, and indirectly, observing the products of DM annihilations (photons, positrons, neutrinos, anti-protons, anti-deuteron) in overdense regions in the galaxy or inside astrophysical objects such as the Sun or the Earth. No positive signal has been found in direct searches except for the two events in CDMS-II~\\cite{Ahmed:2009zw} and the DAMA experiment which has reported a positive signal at 8.2$\\sigma$~\\cite{Bernabei:2010mq}.  Very recently, the CoGeNT collaboration~\\cite{Aalseth:2010vx} has found an indication for excess of events that might be due to scattering of light DM off the nuclei. The first analysis of the XENON100 experiment~\\cite{Aprile:2010um} excludes all positive DM signals, however the interpretation of the data depends on the astrophysical uncertainties~\\cite{McCabe:2010zh} as well as the effective light yield in the low mass region~\\cite{Collar:2010gg,*Collaboration:2010er,*Collar:2010gd}. We will consider these possible signals for DM detection and check whether they can be accommodated in our model. In the case of DAMA, we look into two possible explanations, by either light dark matter or by the inelastic scattering scenario.  As our candidate is a quasi-singlet coupled to the Higgs boson, the interactions which induce the dark matter annihilations at freeze-out are also responsible for elastic or inelastic scattering off nuclei, relevant for direct dark matter searches. In some cases the cross sections of the two processes, annihilations and scattering off nuclei, are related.  All new particles are typically expected to have masses around the electroweak scale and to couple to the SM particles at tree level leading to  rich phenomenology which makes the model testable in the near future. We therefore investigate several experimental bounds including the invisible decay width of the $Z$ into DM pair (if kinematically allowed), the branching ratios of Lepton Flavour Violating (LFV) rare decays and the anomalous magnetic moment of the muon. Furthermore, we comment on the possible signatures at the LHC. It is possible to correlate the flavour structure of the couplings measured at the LHC with the neutrino mass matrix and the data from LFV rare decay searches.  The paper is organised as follows. In \\Secref{sec:model}, the model is presented.  In \\Secref{sec:Scalarmass}, the neutral scalar sector is  analysed. In \\Secref{sec:lepton}, the neutrino mass generation at one loop level as well as effects on  the LFV rare decays and magnetic dipole moment of the muon are discussed. In \\Secref{sec:darkcandidate}, different processes that can give rise to annihilation of dark matter are explored. A discussion of the possibilities of direct and indirect dark matter detection is also included. In \\Secref{sec:lab}, experimental constraints from electroweak precision tests as well as possible collider signatures are studied. Finally, in \\Secref{sec:conclusions}, results are summarised. ",
        "conclusions": "}   In this paper, we have proposed a model that simultaneously explains the missing mass problem of the universe and the tiny neutrino masses. In addition to the SM particle content, there are only a complex scalar singlet and triplet as well as a vector-like electroweak fermionic doublet. We impose an approximate $\\U{1}_{X}$ symmetry which is broken to a remnant $\\mathbb{Z}_2$ symmetry. The unbroken $\\mathbb{Z}_2$ symmetry guarantees the stability of the lightest scalar in the model, $\\delta_1$, a quasi-singlet of $\\SU{2}_L$, which plays the role of dark matter. In the limit of exact $\\U{1}_X$ symmetry, neutrinos are massless and only after the $\\U{1}_X$ is broken to the $\\mathbb{Z}_2$ symmetry, neutrinos acquire a mass term at the one-loop level. The $\\mathbb{Z}_2$ symmetry forbids a tree-level neutrino mass term and the usual seesaw mechanism does not take place. Hence, the smallness of neutrino masses is explained by the small breaking of the $\\U{1}_X$ symmetry as well as the loop suppression. With the minimal particle content of the model, one of the neutrino mass eigenvalues vanishes and the neutrino mass scheme is therefore hierarchical. In order to obtain a non-hierarchical neutrino mass scheme, the minimality of the model has to be relaxed and more vector-like fermionic doublets have to be added. The strongest constraints come from searches for lepton flavour violating processes, in particular $\\mu\\to e\\gamma$, which already probes the relevant parameter space. Future searches for $\\mu\\to e\\gamma$ will provide a very sensitive test of our model.  In this model, DM is produced thermally in the Early Universe. We discussed the different dark matter annihilation channels and identified the dominant one to be the one via Higgs exchange, for $M_1 \\ll m_W$. All other channels are subdominant. The predicted cross section is compatible with the value required to explain the observed DM abundance.  The interactions responsible for DM freeze-out induce also scattering of dark matter off nuclei, relevant for direct DM searches. For our typical values $M_1=70\\GeV$, $\\delta_1$ scatters elastically via Higgs exchange. The obtained scattering cross section is just below the current experimental bound and moderate improvements on the sensitivity can probe part of the relevant parameter space. Our model can also accommodate light dark matter with mass in the few GeV range, which has been invoked to explain the CoGeNT and DAMA results via elastic scattering. This process is mediated by the Higgs exchange and can have the required value for the cross section. For heavier masses, $M_1 \\sim 20$--$50$~GeV, the two events recently reported by CDMS can be interpreted as dark matter elastic scattering with a cross section which is compatible with the predictions of our model. The first results of the XENON100 experiment~\\cite{Aprile:2010um} disfavour most of the parameter region of DAMA, CoGeNT and the two events from CDMS depending on the assumptions on astrophysical uncertainties~\\cite{McCabe:2010zh} and the ratio between electron equivalent energy and nuclear recoil energy $\\mathcal{L}_\\mathrm{eff}$~\\cite{Collar:2010gg,*Collaboration:2010er,*Collar:2010gd}. Further data from the XENON100 experiment as well as other experiments is needed to resolve this uncertainty. We also studied the possibility of inelastic spin independent solution for DAMA. For small mass splittings and small dark matter masses, $\\delta_1$ can scatter inelastically to $\\delta_2$ via $Z$ boson exchange through mixing between scalar singlet and triplet. In our model the mass splitting $\\delta$ can be naturally small due to the $\\U{1}_X$ symmetry. In order to accommodate the solution, the singlet-triplet mixing has to be relatively large and just below the upper bound from  the invisible $Z$ boson decay width. Thus, a slight improvement on the precision of the invisible $Z$ decay width can probe this phenomenologically interesting part of the parameter space in our model.  We demonstrated that bounds on electroweak precision observables do not constrain the model further. Even more, the upper bound on the SM Higgs from electroweak precision data can be relaxed. The new particles can in principle be produced at the LHC and will eventually decay into stable $\\delta_1$ which escapes detection. The second lightest scalar, $\\delta_2$, dominantly decays via $Z$ exchange and might lead to a displaced vertex in the detector for sufficiently large mass splitting $\\delta$, or can decay outside the detector contributing to the missing energy signal. It is possible that  $H \\to \\delta_1\\delta_1$ and $H \\to\\delta_2\\delta_2$ dominate over the SM mode $H\\to b\\bar{b}$, if $\\delta_{1,2}$ are sufficiently light. In this case, the Higgs would decay mainly invisibly. The relevant coupling $\\lambda_L$ for this decay is fixed by the DM annihilation rate. Collider searches for the fermionic doublet can be also performed. By studying the subsequent decay of the charged components of the doublet into a charged lepton, we can determine the Yukawa couplings of these particles to different flavours. The flavour structure of these couplings also determines the flavour structure of the neutrino mass matrix so this provides another method to cross check the model.  Let us finally, comment on an alternative possibility, which leads to a tight connection between neutrino masses and the dark matter abundance~\\cite{Boehm:2006mi}. If the guiding symmetry is not $\\U{1}_X$, but an approximate lepton number $\\U{1}_L$ or $\\U{1}_{B-L}$, the dominant dark matter annihilation channel may be $\\delta_1 \\delta_1 \\to \\nu_\\alpha \\nu_\\beta,\\bar{\\nu}_\\alpha \\bar{\\nu}_\\beta$ resulting in a direct connection between the dark matter abundance and neutrino masses. This leads to an upper bound of the order of $300\\GeV$ on the masses of $E^-_R$, $E'^+_R$, $\\nu_R$ and $\\nu_R^\\prime$ guaranteeing their production at the LHC. In this case, the neutrino flux from dark matter annihilations inside the Sun will be monochromatic with a general flavour composition determined by the flavour structure of the new Yukawa couplings of the model. As recently shown in~\\cite{Esmaili:2009ks}, this can lead to a novel seasonal variation in IceCube which cannot take place in models predicting only a continuous spectrum or democratic neutrino flavour composition.   In summary, we have presented here a model which explains simultaneously the origin of neutrino masses and the dark matter. A global $\\U{1}_X$ symmetry, explicitly broken to a residual $ \\mathbb{Z}_2$ guarantees the smallness of neutrino masses, generated at the loop- level, and the stability of dark matter. The model has a very rich phenomenology, such as lepton flavour violating processes, invisible decays of the $Z$-boson, collider signatures, which will make the model testable in the near future. Dark matter annihilations dominantly proceed via Higgs-exchange. Elastic and/or inelastic scattering off-nuclei can also be induced by the Higgs or $Z$ exchange and can explain the possible signal or hints for dark matter direct detection which have been recently reported. So far, we have considered an explicit breaking of the additional $\\U{1}_X$ symmetry, but a version with a gauged $\\U{1}_X$ symmetry is in preparation."
    },
    "1005/1005.1071_arXiv.txt": {
        "abstract": "{We present a detailed analysis of the far-infrared (-IR) properties of the bright, lensed, $z=2.3$, submillimetre-selected galaxy (SMG), SMM\\,J2135$-$0102 (hereafter SMM\\,J2135), using new observations with {\\it Herschel}, SCUBA-2 and the Very Large Array (VLA). These data allow us to constrain the galaxy's spectral energy distribution (SED) and show that it has an intrinsic rest-frame 8--1000-$\\mu$m luminosity, $L_{\\rm bol}$, of $(2.3\\pm0.2)\\times 10^{12}$\\,L$_{\\odot}$ and a likely star-formation rate (SFR) of $\\sim 400$\\,M$_{\\odot}$\\,yr$^{-1}$. The galaxy sits on the far-IR/radio correlation for far-IR-selected galaxies. At $\\gs 70$\\,$\\mu$m, the SED can be described adequately by dust components with dust temperatures, \\td\\ $\\sim$ 30 and 60\\,{\\sc k}.  Using SPIRE's Fourier Transform Spectrometer (FTS) we report a detection of the \\cii\\,158$\\mu$m cooling line. If the \\cii, CO and far-IR continuum arise in photo-dissociation regions (PDRs), we derive a characteristic gas density, $n\\sim 10^{3}$\\,cm$^{-3}$, and a far-ultraviolet (-UV) radiation field, $G_0$, $10^3\\times$ stronger than the Milky Way. $L_{\\rm [CII]}/L_{\\rm bol}$ is significantly higher than in local ultra-luminous IR galaxies (ULIRGs) but similar to the values found in local star-forming galaxies and starburst nuclei.  This is consistent with SMM\\,J2135 being powered by starburst clumps distributed across $\\sim$2\\,kpc, evidence that SMGs are not simply scaled-up ULIRGs. Our results show that SPIRE's FTS has the ability to measure the redshifts of distant, obscured galaxies via the blind detection of atomic cooling lines, but it will not be competitive with ground-based CO-line searches. It will, however, allow detailed study of the integrated properties of high-redshift galaxies, as well as the chemistry of their interstellar medium (ISM), once more suitably bright candidates have been found.} ",
        "introduction": "Submillimetre (submm) surveys have uncovered a population of intrinsically luminous, but highly obscured, galaxies at high redshift.  However, even with instrinsic luminosities of $\\sim10^{13}$\\,L$_{\\odot}$ \\citep[e.g.][]{ivison98}, the brightest SMGs are still challenging targets for observational studies. In the submm and far-IR, where the bulk of their luminosity escapes, the brightest SMGs have observed flux densities of only $\\sim10$\\,mJy at 850\\,$\\mu$m, peaking at $\\sim50$\\,mJy at the wavelengths probed by {\\it Herschel}. To alleviate this photon starvation, submm surveys often exploit gravitational lensing via massive, foreground galaxy clusters, thereby enhancing the apparent brightness of SMGs at all wavelengths \\citep[e.g.][]{smail97, chapman02, cowie02}.  Recently, \\citet{swinbank10a} exploited the cluster lensing technique using the Large Apex Bolometer Camera \\citep[LABOCA --][]{siringo09} on the 12-m Atacama Pathfinder Experiment (APEX) telescope to map the cluster, MACS\\,J2135$-$01 ($z=0.325$), and thereby discovered SMM\\,J2135, an SMG with $S_{\\rm 870\\mu m} =106$\\,mJy. Its brightness is due to very high amplification (by $32.5\\pm4.5$) by the foreground cluster (similarly bright sources may have recently been unearthed by the South Pole Telescope -- \\citealt{vieira10}).  The lens model for SMM\\,J2135 is well constrained and its redshift ($z=2.3259\\pm 0.0001$, derived from the detection of CO $J$=1--0 in a blind search) and intrinsic flux ($3.3\\pm0.5$\\,mJy) are typical of SMGs found close to the confusion limit in submm surveys. SMM\\,J2135 thus presents an opportunity to study a member of this important population at high signal-to-noise and with the spatial and spectral resolution necessary to determine the detailed far-IR spectral properties of SMGs. Due to the high magnification, it is feasible to apply some of the observational tools used on local star-forming galaxies to understand the processes of star formation at high redshift.  Indeed, we can employ diagnostics capable of determining the flux of ionising radiation and the SFR, thus determining the state of the overwhelming majority of the atomic and molecular gas in this galaxy \\citep{wolfire90, hollenbach99, kaufman99}.  In this paper we present spectroscopic and photometric far-IR/submm measurements of SMM\\,J2135 made using {\\it Herschel} \\citep{pilbratt10}.  We also include new observations with the James Clerk Maxwell Telescope (JCMT) and VLA. We use these observations to constrain the SED of SMM\\,J2135 and measure or set firm limits for the line fluxes from the main atomic cooling lines. ",
        "conclusions": "We have delineated the far-IR SED of a highly magnified (but intrinsically typical) SMG, SMM\\,J2135, at $z=2.3$. Its rest-frame 8--1000-$\\mu$m and 1.4-GHz luminosities are $2.3\\times 10^{12}$\\,L$_{\\odot}$ and $9\\times 10^{23}$\\,W\\,Hz$^{-1}$, with SFR $\\sim$ 400\\,M$_{\\odot}$\\,yr$^{-1}$, and it sits on the far-IR/radio correlation for starburst galaxies.  {\\it Herschel} FTS spectroscopy detects the redshifted [C\\,{\\sc ii}]\\,158\\,$\\mu$m emission line, allowing us to investigate the properties of its ISM. The line luminosity suggests that the mass of \\cii\\ is $\\sim$25\\% of the molecular gas, similar to the ratio found in local starbursts.  We use CO(1--0), \\cii\\ and $L_{\\rm bol}$ to investigate the ISM's physical conditions.  From a comparison with PDR models, we derive a far-UV radiation field, $G_0$, which is $\\sim10^3\\times$ higher than that in the Milky Way, but comparable to those found in ULIRGs.  In contrast, we find a characteristic density, $n\\sim 10^{3}$\\,cm$^{-3}$, which is lower than seen in ULIRGs, but comparable to values seen in local star-forming galaxies and nuclei, as well as a small number of high-redshift systems where similar measurements have been made. Together these results suggest that SMM\\,J2135 has a SFR intensity similar to that seen in local ULIRGs, but distributed over a larger volume. This is consistent with the $\\sim$2-kpc distribution of star formation across this galaxy \\citep[][]{swinbank10a} and previous suggestions of extended star formation in SMGs \\citep[e.g.][]{biggs08}.  Our results show that SPIRE's FTS has the ability to measure the redshifts of suitably bright and distant, obscured galaxies via detection of atomic cooling lines such as \\cii. However, we estimate that $\\gs$10-hr integrations will be required and this is not competitive with blind, ground-based CO-line searches \\citep[e.g.][]{weiss09}, as evidenced by the ease with which the redshift of SMM\\,J2135 was determined using Zpectrometer on the Green Bank Telescope \\citep{swinbank10a}.  Nevertheless, our results show that facilities such as {\\it Herschel} and SCUBA-2 will allow detailed study of the integrated properties of high-redshift galaxies (through SED modelling), as well as the chemistry of their ISM."
    },
    "1005/1005.1592_arXiv.txt": {
        "abstract": "{Comet C/2006 W3 (Christensen) was observed in November 2009 at 3.3 AU from the Sun with {\\it Herschel}.The PACS instrument acquired images of the dust coma in 70-$\\mu$m and 160-$\\mu$m filters, and spectra covering several H$_2$O rotational lines. Spectra in the range 450--1550 GHz were acquired with SPIRE. The comet emission continuum from 70 to 672 $\\mu$m was measured, but no lines were detected. The spectral energy distribution indicates thermal emission from large particles and provides a measure of the size distribution index and dust production rate. The upper limit to the water production rate is compared to the production rates of other species (CO, CH$_3$OH, HCN, H$_2$S, OH) measured with the IRAM 30-m and Nan\\c{c}ay telescopes. The coma is found to be strongly enriched in species more volatile than water, in comparison to comets observed closer to the Sun. The CO to H$_2$O production rate ratio exceeds 220\\%. The dust to gas production rate ratio is on the order of 1.  } ",
        "introduction": "Direct imaging shows that distant activity is a general property of cometary nuclei \\citep[e.g.,][]{mazo09}. It is attributed to the sublimation of hypervolatile ices, such as CO or CO$_2$, or to the release of volatile species trapped in amorphous water ice during the amorphous-to-crystalline phase transition \\citep{pria04}. Indeed, at heliocentric distances $r_h$ larger than 3--4 AU, the sublimation of water, the major volatile in cometary nuclei, is inefficient. Characterizing the processes responsible for distant activity is important for understanding the structure and composition of cometary nuclei, their thermal properties and their evolution upon solar heating. However, detailed investigations of distant nuclei are sparse. The best studied objects are the distant comet 29P/Schwassmann-Wachmann 1, where CO, CO$^+$, and CN were detected at 6 AU from the Sun, and C/1995 O1 (Hale-Bopp), whose exceptional activity allowed us to detect several molecules and radicals farther than 3 AU --- including CO up to 14 AU \\citep{biv02,rauer03}.  Comet C/2006 W3 (Christensen) was discovered in November 2006 at $r_h$ = 8.6 AU from the Sun with a total visual magnitude $m_v$ $\\sim$18. This long-period comet passed perihelion on 9 July 2009 at $r_h$ = 3.13 AU. Because of its significant brightness ($m_v$ $\\sim$ 8.5 at perihelion), it was an interesting target for the study of distant cometary activity. We report here on observations undertaken at $r_h$ = 3.3 AU post-perihelion with the PACS \\citep{Poglitsch10} and SPIRE \\citep{Griffin10} instruments on {\\it Herschel} \\citep{Pilbratt10}, in the framework of the {\\it Herschel} Guaranteed Time Key Project called ''Water and related chemistry in the Solar System'' \\citep{hart09}. These observations are complemented by production rate measurements of several species using the Nan\\c{c}ay radio telescope and the 30-m telescope of Institut de Radioastronomie millim\\'etrique (IRAM) at 3.2--3.3 AU pre- and post-perihelion. ",
        "conclusions": "Comet Christensen was a distant comet. Nevertheless, the continuum was clearly detected by PACS and SPIRE, providing useful constraints on the properties of the cometary dust. Although water emission was not detected in this object, the limits obtained are significant. The prospects for future cometary studies with {\\it Herschel} are thus very good."
    },
    "1005/1005.1889_arXiv.txt": {
        "abstract": "{We use {\\it Herschel} Space Observatory data to place observational constraints on the peak and Rayleigh-Jeans slope of dust emission observed at 70-500~$\\mu$m in the nearby spiral galaxy \\object{M81}.  We find that the ratios of wave bands between 160 and 500~$\\mu$m are primarily dependent on radius but that the ratio of 70 to 160~$\\mu$m emission shows no clear dependence on surface brightness or radius.  These results along with analyses of the spectral energy distributions imply that the 160-500~$\\mu$m emission traces 15-30~K dust heated by evolved stars in the bulge and disc whereas the 70~$\\mu$m emission includes dust heated by the active galactic nucleus and young stars in star forming regions.} ",
        "introduction": "The {\\it Herschel} Space Observatory\\footnote{{\\it Herschel} is an ESA space observatory with science instruments provided by Principal Investigator consortia. It is open for proposals for observing time from the worldwide astronomical community. } \\citep{pilbrattetal10} provides an unprecedented view of the far-infrared and submillimetre emission from nearby galaxies.  At wavelengths of 70-160 $\\mu$m, the PACS instrument \\citep{poglitschetal10} can produce images with resolutions of $6^{\\prime\\prime}$-$12^{\\prime\\prime}$ that are superior to what can be achieved with the {\\it Spitzer} Space Telescope.  At 250-500~$\\mu$m, the SPIRE instrument \\citep{getal10} produces images with unprecedented sensitivities to diffuse and point-like submillimetre emission.  We can use these data to construct spectral energy distributions (SEDs) that sample the peak and Rayleigh-Jeans side of thermal dust emission, thus allowing us to probe the coldest dust components in nearby galaxies and place superior constraints on dust temperatures and masses.  As part of the Very Nearby Galaxies Survey (VNGS), we have imaged the spiral galaxy M81 (NGC 3031) at 70, 160, 250, 350, and 500~$\\mu$m with PACS and SPIRE.  M81 is a nearby \\citep[$3.63\\pm0.13$~Mpc;][]{fetal01} SA(s)ab \\citep{ddcbpf91} galaxy at an inclination of $59.0^\\circ$ \\citep{detal08} with well defined spiral arms.  The {\\it Herschel} data allow us to extract SEDs for $\\sim0.7$~kpc subregions that are small enough that we can distinguish arm and interarm regions within M81. We use these data to explore the dust temperatures and masses and to understand the heating sources for the dust. ",
        "conclusions": ""
    },
    "1005/1005.3242_arXiv.txt": {
        "abstract": "We have measured the physical properties of polar coronal holes from the minimum activity phase of solar cycle 23 (1996--1997) to the present minimum of solar cycle 24 (2007--2009) using the UVCS instrument on SOHO. Observations in H~I Lyman alpha (121.6 nm) and O~VI (103.2, 103.7 nm) provide spectroscopic diagnostics of proton and O$^{5+}$ bulk outflow velocities and velocity distributions as a function of heliocentric distance above the poles of the Sun. These observations have allowed us to follow the changes in the physical properties of the polar coronal holes during solar cycle 23 and its approach to the current minimum. Recent ground- and space-based observations have reported a variety of phenomena associated with the current minimum. We present the comparison of observed oxygen line intensities, line ratios, and profiles for polar coronal holes at both minima and during solar cycle 23 and show how this new minimum manifests itself in the ultraviolet corona. The comparison of the physical properties of these two minima as seen by UVCS in the extended corona, now possible for the first time, may provide crucial empirical constraints on models of extended coronal heating and acceleration for the fast solar wind. ",
        "introduction": "Polar coronal holes are magnetically open large-scale spatial structures that are present for most of the solar cycle, and their role is far from being well understood. It is clear that strong connections exist between large coronal holes and the highest-speed wind streams (Krieger et al.\\  1973; Zirker 1977).  Recent ground- and space-based observations have reported a variety of phenomena associated with the current solar minimum. Solar cycle 23 was a longer than expected activity cycle, and the current solar cycle 24 minimum did not produce a quiescent equatorial streamer belt. The solar wind during the previous sunspot minimum presented a stable bimodal structure, composed of low-speed streams around the equator and high-speed streams over the poles. However, the solar wind from the current minimum period differs from this typical configuration. During this minimum, large low-latitude coronal holes, which are also sources of fast solar wind, were located at the equator for an extended period of time, unlike in the previous minimum (see, e.g., Galvin et al.\\  2008; Miralles 2008; Tokumaru et al.\\  2009).  The current polar coronal holes and the fast solar wind have significantly different properties than at the 1996--1997 minimum. {\\em In situ} observations of the solar wind from both polar coronal holes show that the fast solar wind is slightly slower, less dense, cooler, and has a lower momentum flux than during the 1996--1997 solar minimum (McComas et al.\\  2008; Issautier et al.\\  2008). Magnetic field measurements from the Wilcox Solar Observatory and the Michelson Doppler Imager (MDI) (Scherrer et al.\\  1995) on the {\\em{Solar and Heliospheric Observatory}} ({\\em{SOHO}}) show that the polar fields are weaker than in the preceding minimum (Sun et al.\\  2008). In addition, polar hole perimeter measurements using images from the Extreme ultraviolet Imaging Telescope (EIT) (Delaboudini\\`{e}re et al.\\  1995) on {\\em SOHO} indicate a reduction of the coronal hole area in both poles of about 15\\% between 1996 and 2007 (Kirk et al.\\  2009).  The exact manner in which the plasma in coronal holes is heated and accelerated is still unresolved. An improvement in our understanding of the physical processes responsible for the solar wind have come from the past decade of observations, analysis, and theoretical work associated with the {\\em SOHO} mission. For example, the importance of magnetohydrodynamic (MHD) waves has been emphasized by the observations of the Ultraviolet Coronagraph Spectrometer (UVCS) (Kohl et al.\\  1995, 2006) on {\\em SOHO} that heavy ions are heated to hundreds of times the temperatures of protons and electrons, and that the velocity distributions are anisotropic, indicating Alfv\\'{e}n wave dissipation via ion cyclotron resonance (Kohl et al.\\  1997, 1998, 2006; Cranmer et al.\\  1999, 2008). These advances were based on observations of polar coronal holes at the last solar minimum (1996--1997). Coronal holes observed since that time have higher electron densities, lower kinetic temperatures, slower outflow velocities, and less divergent magnetic field geometries (Miralles et al.\\  2001a,b, 2002, 2004; Miralles 2006). It appears that the physical processes controlling the extended heating and acceleration may depend on the density and on the magnetic field geometry.  This paper provides an overview of the physical properties of the polar coronal holes during solar cycle 23 and the current cycle 23/24 minimum derived from observations by UVCS. In particular, spectroscopic diagnostics in polar coronal holes are discussed.  The resulting plasma properties of these coronal holes are compared. For a review of the physics of coronal heating and solar wind acceleration of the associated high-speed solar wind, see Cranmer et al.\\  (2010, these proceedings). ",
        "conclusions": "UVCS/{\\em{SOHO}} spectroscopic observations of polar coronal holes during solar cycle 23 have shown marked variations of ion properties in the solar wind acceleration region. Polar coronal holes seem to exhibit different kinetic temperatures and acceleration rates. These observations show that the north and south polar coronal holes in 2007--2009 exhibited similar extreme properties (intensities and line widths) as those measured by UVCS in the 1996--1997 previous minimum. From the analysis of the line widths, we can infer that the O~VI heating seems as strong as in the previous solar minimum. A slight decrease in the polar coronal O~VI intensities is seen that may be consistent with {\\em in situ} measurements of lower electron temperatures and densities in the fast solar wind. The coronal magnetic field is $\\sim$ 40\\% lower (Sheeley 2010, these proceedings), but this magnitude of decrease in the ion heating may not be present in the corona."
    },
    "1005/1005.2134.txt": {
        "abstract": "We extend the  formalism of weak interaction processes, obtaining new  expressions for the transition rates,  which  greatly facilitate numerical calculations, both for neutrino-nucleus reactions and muon capture. Explicit violation of CVC hypothesis by the Coulomb field, as well as development of a sum rule approach for the inclusive cross sections have been worked out. We have done a thorough  study of exclusive (ground state) properties of  $^{12}$B and $^{12}$N  within  the projected quasiparticle random phase approximation (PQRPA). Good agreement with experimental data achieved  in this way put in evidence the limitations of standard  RPA and the QRPA models, which come from the inability of the RPA in opening the $p_{3/2}$ shell, and from the non-conservation of the number of particles in the QRPA. The inclusive neutrino/antineutrino ($\\nu/\\tilde{\\nu}$) reactions $^{12}$C($\\nu,e^-)^{12}$N and $^{12}$C($\\tilde{\\nu},e^+)^{12}$B are calculated within both the PQRPA,  and the relativistic QRPA (RQRPA). It is found that the magnitudes of the resulting cross-sections: i) are  close to the sum-rule limit at low energy, but significantly smaller than this limit at high energies both for $\\nu$ and $\\tilde{\\nu}$, ii) they steadily increase when the  size of the configuration space is augmented, and particulary for $\\nu/\\tilde{\\nu}$ energies $> 200$ MeV, and iii) converge  for sufficiently large configuration space and final state spin. The quasi-elastic   $^{12}$C($\\nu,\\mu^-)^{12}$N cross section   recently measured in the MiniBooNE experiment is briefly discussed. We study the decomposition of the inclusive cross-section based on the degree of forbiddenness  of different multipoles. A few words are dedicated to the $\\nu/\\tilde{\\nu}$-$^{12}$C charge-exchange reactions related with astrophysical applications. ",
        "introduction": "\\label{Sec1} The massiveness of neutrinos and the related oscillations are strongly sustained by many experimental works involving atmospheric, solar, reactor and accelerator neutrinos ~\\cite{Ath96,Ath98,Agu01,Fuk98,Aha05,Ara04,Ahn03}. The  subsequent experimental goal   is to determine precisely the various parameters of the Pontecorvo-Maki-Nakagawa-Sakata (PMNS) neutrino mass matrix, absolute masses of different flavors of neutrinos, CP violation in neutrino sector, etc. To address  these problems several analyses of neutrino oscillation data are presently going on. At the same time, several experiments  presently collect data, and others are planned. Accelerator experiments, experiments with neutrinos from $\\nu$-factories, $\\beta$-beams, etc., are also planned and designed, as well as some experiments with natural $\\nu$-sources like solar neutrinos, atmospheric neutrinos, or antineutrinos from nuclear reactors.   The neutrino-nucleus scattering on $^{12}$C is important because this nucleus is a component of many liquid scintillator detectors. Experiments such as LSND~\\cite{Ath96,Ath98}, KARMEN~\\cite{Mas98,Arm02}, and LAMPF~\\cite{All90,Kra92}  have used $^{12}$C to search for neutrino oscillations, and for measuring neutrino-nucleus cross sections. %{\\color {red} Present atmospheric and accelerator based neutrino oscillation experiments  also involve $^{12}$C, and operate at neutrino energies $E_\\nu\\sim 1$ GeV in order to access the relevant regions of oscillation parameter space. This is the case of the SciBar detector~\\cite{SciBar}, where the molecule $C_8H_8$ is involved, and the  MiniBooNE detector ~\\cite{MiniBooNE}, which uses the light mineral oil containing the molecule $CH_2$. The $^{12}$C target will be used in several planned experiments, such as the spallation neutron source (SNS) at Oak Ridge National Laboratory (ORNL)~\\cite{Efr05}, and the LVD (Large Volume Detector) experiment~\\cite{Aga07}, developed by the INFN in Gran Sasso.  %\\newpage For the planned experimental searches  of  supernovae neutrino signals, which involve $^{12}$C as scintillator liquid detector,  the precise knowledge of neutrino cross sections  of $^{12}$N and $^{12}$B ground-states, \\ie of $\\sigma_{e^-}(E_\\nu,1^+_{1})$, and $\\sigma_{e^+}(E_{\\tilde{\\nu}},1^+_{1})$  is very important. In fact, in the LVD experiment~\\cite{Aga07} the number of events detected during the supernova explosion are estimated  by convoluting  the neutrino supernova flux with: i)  the interaction cross sections, ii) the efficiency of the detector, and iii) the number of target nuclei. For the carbon content of the LVD detector have been used so far $\\sigma_{e^-}(E_\\nu,1^+_{1})$, and $\\sigma_{e^+}(E_{\\tilde{\\nu}},1^+_{1})$, as obtained from the Elementary Particle Treatment (EPT)~\\cite{Fuk88}. %{\\color{blue} Moreover, as an update of the LVD experiment related to supernovae neutrinos detection (where $^{12}$C will also be employed), there is ongoing design study  concerning large size scintillator detectors, called LAGUNA, where a 50 kt scintillator LENA is being considered \\cite{Aut07}. %} %}  On the other hand,  as the $^{12}$C nucleus forms one  of the onion-like shells of a large star before collapse, it is also important for astrophysics studies. Concomitantly, several authors~\\cite{Lun03,Dig03,Aga07,Dua07,Das08,Das10,Mez10} have recently stressed the importance of  measuring supernova neutrino oscillations. They claim that a supernova explosion represents a unique scenario for further study of the PMNS matrix. The corresponding  neutrinos, which carry  all flavors were observed in only one occasion (SN1987A), have an energy $E_\\nu \\lesssim 100$~MeV \\cite{Str06}, and are also studied through the interactions with carbon nuclei in the liquid scintillator.  %{\\color {red} Thus, the main interest in the neutrino/antineutrino-$^{12}$C charge-exchange cross sections comes from the neutrino oscillations, and precise knowledge of the cross sections in the neutrino energies  going from a  few  MeV\u00b4s up to a few GeV\u00b4s is required. Up to quite recently the only available  experimental information on reactions was that for the flux-averaged cross-sections: i) $^{12}$C$(\\nu_e,e^-)^{12}$N in the DAR region: $E_{\\nu_e} < 60$~MeV~\\cite{Ath97,Aue01,Zei98}, and ii) $^{12}$C$(\\nu_\\mu,\\mu^-)^{12}$N in the  DIF region: $127$ MeV $\\leq E_{\\nu_\\mu}\\leq 300$ MeV~\\cite{Ath97a,Aue02a,LSND}. In last few years, however, several experimental programs at MiniBooNE~\\cite{MiniBooNE Collaboration}, K2K~\\cite{K2K Collaboration}, and  SciBooNE~\\cite{SciBooNE Collaboration} yield results on the ($\\nu_\\mu,^{12}$C)  cross section for $0.4$ GeV $\\leq E_{\\nu_\\mu}\\leq 1.7$ GeV. It is well known  that for $E_\\nu$ larger than a few hundreds  MeV's, besides the quasi-elastic (QE) channel, many inelastic channels are open and pion production becomes important. %\\cite{Lei09,Mart09,Nak10,Lei10,Saj10}. In fact, there have  been quite active experimental efforts to investigate neutrino-induced coherent single-pion production in the  $\\Delta$-excitation region of $^{12}$C. Starting approximately at the  threshold coming from the pion and charged lepton masses ($\\mass_\\pi$ and $\\mass_\\ell$), the $\\pi+\\ell$ %pion plus lepton production cross section steadily increases with the neutrino energy becoming larger than the quasi-elastic one for $E_\\nu\\lsim 1.5$ GeV~\\cite{MiniBooNE Collaboration, K2K Collaboration, SciBooNE Collaboration}.   From the  theoretical side there have been great efforts to understand the  nuclear structure within the triad $\\{^{12}$B,$^{12}$C,$^{12}$N$\\}$. In the seminal   work of O'Connell,  Donelly,  and Walecka ~\\cite{Con72} a unified analysis of electromagnetic, and semileptonic weak interactions was presented. To describe the nuclear dynamics they have used the particle-hole Tamm-Dancof Approximation (TDA) within a very small single-particle space \\footnote{From now on a    single-particle (s.p.) space that includes all orbitals  within  $N$ harmonic oscillator (HO) shells will be  labeled as space $S_{N}$.} ($S_2$ $\\equiv \\{1s_{1/2}$, $1p_{3/2}$, $1p_{1/2}$, $1d_{5/2}$, $2s_{1/2}\\}$)~\\cite{Don70}. To achieve  agreement with experiments for the $\\beta^\\pm$-decays, and $\\mu$-capture they were forced to use an overall reduction factor $\\xi^2$ of the order of $4$ ($2$) for even (odd) parity states. They have also pointed out that this factor would become totally unnecessary with  use of a better nuclear model able to open the $1p_{3/2}$ shell.    Rather thorough comparisons of $2s1d$ and $2p1f$ shell-model predictions with measured allowed $\\beta$-decay rates have yielded a simple, phenomenological effective axial coupling $g_{\\sss A}^{\\sss}=1$ that should be used rather than the bare value \\cite{Bro85,Cas87,Ost92,Mar96}. This observation is the basis for many nuclear model estimates of the  Gamow-Teller (GT) response that governs allowed neutrino cross sections. In Ref.~\\cite{Con72} $g_{\\sss A}^{\\sss}=1.23$ was used based on a study of neutron $\\beta$-decay, and,  as the analyzed  processes were dominantly of the axial-vector type, the use of $g_{\\sss A}^{\\sss}=1$ would  have diminished the reduction factors $\\xi^2$ in an appreciable way.    In the Random Phase Approximation (RPA),  besides the TDA forward-going amplitudes,  the backward-going amplitudes are present as well. However, these additional  RPA amplitudes did not help to open the $1p_{3/2}$ shell in the continuum RPA (CRPA) calculations of Kolbe, Langanke, and Krewald~\\cite{Kol94}. Thus,  as in the case of the TDA used  in Ref.~\\cite{Con72}, to get agreement with data for the ground state triplet $T=1$ ($\\beta^\\pm$-decays, $\\mu$-capture, and the exclusive $^{12}$C$(\\nu_e,e^-)^{12}$N reaction) their calculations were  rescaled  by a factor $\\cong 4$.  The main aim of the CRPA is to  describe appropriately not only the bound states but also  the virtual (quasi-bound), resonant, and  continuum states, which are treated as bound states in the RPA. However, this superiority has not been evidenced so far in numerical calculations.  For instance, in the case of $\\mu$-capture rates in $^{16}$N  the two methods agree with each other quite well for the  $0^{-}$ and $1^{-}$  states, while the RPA result is preferred for the $2^{-}$  state ~\\cite{Kol94a}.  To open the $1p_{3/2}$ shell one has to introduce pairing correlations. This is done within the Shell Model (SM)~\\cite{Hay00,Vol00,Suz06}, which reproduces quite well both i) the experimental flux-averaged exclusive, and inclusive cross sections for the $^{12} \\rm C(\\nu_e,e^-)^{12}\\rm N$ DAR~\\cite{Ath97,Aue01,Zei98}, and $^{12}\\rm C(\\nu_\\mu,\\mu^-)^{12}\\rm N$ DIF~\\cite{Ath97a} reactions, and ii) the $\\mu^- + {^{12}\\rm C}\\rightarrow \\nu_{\\mu}+{^{12}\\rm B}$ muon-capture modes~\\cite{Mil72,Mea01,Sto02}.  The quasiparticle RPA (QRPA) also opens the $1p_{3/2}$ shell by means of the pairing interaction. However,  it  fails as well in accounting for the exclusive processes to the isospin triplet $T=1$ in $^{12}\\rm C$, because a new problem emerges, as  first observed by Volpe \\etal~\\cite{Vol00}.   They noted that within the QRPA the lowest state in ${^{12}\\rm N}$ irremediably turned out not to be the most collective one.  Later it was shown \\cite{Krm02,Krm05,Sam06} that: 1) the origin of this difficulty arises from the degeneracy among  the four lowest proton-neutron two-quasiparticle ($2qp$) states $\\ket{1p_{1/2}1p_{3/2}}$, $\\ket{1p_{3/2}1p_{3/2}}$, $\\ket{1p_{1/2}1p_{1/2}}$ and $\\ket{1p_{3/2}1p_{1/2}}$, which, in turn, comes from the fact that for $N=Z=6$ the quasiparticle energies $E_{1p_{1/2}}$ and $E_{1p_{3/2}}$ are very close to  each other, and 2) it is imperative to use the projected QRPA (PQRPA) for a physically sound description of  the weak processes among  the ground states of the triad $\\{{{^{12}\\rm B},{^{12}\\rm C},{^{12}\\rm N}}\\}$~\\cite{Krm02,Krm05,Sam06}; see  Figs. 2 and 3 in Ref.~\\cite{Krm05}.  In summary, neither the CRPA nor the QRPA are the appropriate nuclear models to describe the ``fine structure\" of exclusive  charge-exchange processes around $^{12}$C, and they only can be used for global inclusive descriptions. Of course, the same is valid for the relativistic RPA (RQRPA) that  has recently been applied with success in calculations of inclusive charged-current neutrino-nucleus reactions in $^{12}$C, $^{16}$O, $^{56}$Fe, and  $^{208}$Pb~\\cite{Paa07}, and total muon capture rates on a large set of nuclei from $^{12}$C to $^{244}$Pu~\\cite{Mar09}. The continuum QRPA (CQRPA) would have to be superior to the QRPA for the same reasons that the CRPA  would have to be better than the RPA. Nevertheless, neither this superiority  has been put in evidence by  numerical calculations~\\cite{Hag01,Rod08}. Finally, it is clear that the nuclear structure descriptions  inspired on the Relativistic Fermi Gas Model (RFGM)~\\cite{Smi72,Nie04,Val06}, which do not involve multipole expansions,  should only  be used for inclusive quantities.  When the effects due to resonant and continuum states are considered, as it is done within the CRPA and CQRPA, the  spreading in strength of the hole states in the inner shells should also be taken into account for the sake of consistency. In fact, a single-particle state $j$ that is  deeply bound in the parent nucleus,  after a weak interacting process  can become a highly excited hole-state $j^{-1}$ in the continuum of the residual nucleus. There it is suddenly mixed with more complicated configurations (2h1p, 3h2p, \\dots excitations, collective states, and so on) spreading its strength in a relatively wide energy interval~\\cite{Ma85}\\footnote{One should keep in mind that the mean life of $^{12}$N and $^{12}$B are, respectively, $11.0$ and $20.2$  ms, while  strong interaction times are of the order of $10^{-21}$ s.}. This happens, for instance, with the $1s_{1/2}$ orbital in $^{12}$C, that is separated  from the $1p_{3/2}$ state by approximately $23$ MeV, which is enough to break the $12$ particle system, where the energy of the last excited state amounts to $ 11.5$ MeV in $^{12}$N, and $ 16.5$ MeV in $^{12}$B  channels. Although  the detailed structure and fragmentation of hole states are still not well known, the exclusive knockout reactions provide a wealth of information on the structure of single-nucleon states of nuclei. Excitation energies  and widths of proton-hole states were systematically measured with quasifree (p, 2p) and $(e, e' $p) reactions, which revealed the existence of inner orbital shells in nuclei~\\cite{Ja73,Fr84,Be85,Le94,Ya96,Ya01,Yo03,Ya04,Ko06}.  In the  TDA calculation of Ref.~\\cite{Con72} the  $S_2$ space has been used, which  extends only from $13.77$ MeV up to $30.05$ MeV, embracing, respectively, $1, 2, 2, 1$, and $1$ negative parity states $J^\\pi=0^-, 1^-, 2^-, 3^-$, and $4^-$, and $1, 2, 2$, and $1$ positive parity states $J^\\pi=0^+, 1^+, 2^+$, and $3^+$. With such  small configuration space, the neutrino cross sections $\\sigma_e(E_\\nu)$, and  $\\sigma_\\mu(E_\\nu)$  have been evaluated up to a neutrino energy $E_\\nu$ %\\red{somewhat} larger than of $0.6$ GeV, and extrapolated up to $20$ GeV. In recent years, however,  large configuration spaces have been used in the evaluation of  QE cross sections for $E_\\nu\\sim 1$ GeV. For instance,  Amaro \\etal~\\cite{Ama05} have employed the single-particle SM (TDA without the residual interaction) in a semirelativistic description of quasielastic neutrino reactions $(\\nu_\\mu,\\mu^-)$ on $^{12}$C going up to $E_\\nu= 1.5$ GeV, and including multipoles $J^\\pi\\leq 47^{\\pm}$. Good agreement with the RFGM was obtained for several choices of kinematics of interest for the ongoing neutrino oscillation experiments. Kolbe \\etal~\\cite{Kol03} have also achieved an excellent agreement between the RFGM and the CRPA calculations of the total cross section and the angular distribution of the  outgoing electrons in $^{16}{\\rm O}(\\nu_e,e)X$ for $E_\\nu\\le  0.5$ GeV. They have considered  states  up to $J^\\pi=9^{\\pm}$ only,  and didn't specify the configuration space used. Moreover, Valverde \\etal~\\cite{Val06} have analyzed the theoretical uncertainties of the RFGM developed in~\\cite{Nie04}  for the $(\\nu_e,e^-)$, and $(\\nu_\\mu,\\mu^-)$ cross sections in  $^{12}$C, $^{16}$O, and $^{40}$Ca for $E_\\nu\\le  0.5$ GeV. The work of Kim \\etal~\\cite{Kim08} should also be mentioned where were studied the effects of strangeness on the $(\\nu_\\mu,\\mu^-)$ and  $({\\tilde\\nu}_\\mu,\\mu^+)$ cross sections in  $^{12}$C for incident energies between  $0.5$ MeV and $1.0$ GeV,  within the framework of a relativistic single-particle model. Quite recently, Butkevich~\\cite{But10} has also studied the scattering of muon neutrino on carbon targets for neutrino energies up to $2.8$ GeV within  a relativistic shell-model approach without specifying the model space. %{\\color{red}   For relatively large neutrino energies ($E_{\\nu_e}\\gsim \\mass_\\pi$, and $E_{\\nu_\\mu}\\gsim \\mass_\\pi+\\mass_\\mu$) to the above-mentioned QE cross sections  should be added the pion production cross section,  as done, for instance, in Refs.~\\cite{Mart09,Lei09}. One should also note that $\\sigma_e(E_{\\nu_e})$ and $\\sigma_\\mu(E_{\\nu_\\mu})$ coincide with each other asymptotically. This is clearly put in evidence in  the Extreme Relativistic Limit (ERL) where $|\\pb_\\ell|/E_\\ell\\go 1$, and the neutrino-nucleus cross sections depend on $\\mass_\\ell$ only trough the threshold energy, as can be seen from the Appendix~\\ref{C} of the present work. The  Figure 4 from Ref.~\\cite{Kur90} is also illustrative in this respect.    Therefore, we  focus our attention only on the quasi-elastic cross section $\\sigma_e(E_{\\nu_e})$ since, at muon-neutrino energies involved in the MiniBooNE experiment~\\cite{MiniBooNE}, it is equal to $\\sigma_\\mu(E_{\\nu_\\mu})$   for all practical purposes.  One of main objectives in the present study is to analyze the effect of the size of the configuration space up to neutrino energies of several hundred MeV. As in several previous works \\cite{Con72,Mar96,Kol94,Kol94a,Hay00,Vol00,Suz06,Krm02,Krm05,Sam06,Paa07,Mar09, Ama05,Kol03,Val06,Kim08,Kur90} this will be done in first order perturbation theory. The consequences  of the particle-particle force in the S = 1, T = 0 channel, within the PQRPA will also  be examined.  The importance of this piece of the residual interaction was discovered more than 20 years ago by  Vogel and Zirnbauer~\\cite{Vog86} and Cha~\\cite{Cha87}, and since then the QRPA  became  the most frequently used nuclear structure method for evaluating double  $\\beta$-decay rates.  A few words will be devoted as well as to the non-relativistic formalisms for neutrino-nucleus scattering. The most popular  one was developed by  the Walecka group~\\cite{Con72,Don79,Hax79,Wal95}, where the nuclear transition matrix elements are classified as Coulomb, longitudinal, transverse electric, and transverse magnetic multipole moments. We feel that these denominations might be convenient when  discussing simultaneously  charge-conserving, and charge-exchange processes, but seems  unnatural when one considers  only the last ones. %{\\it As a matter of fact, this terminology is not often used in  nuclear $\\beta$-decay,  $\\mu$-capture, and charge-exchange reactions where  one only speaks   of vector and axial matrix elements with different degrees of   forbiddenness: allowed (GT and Fermi), first forbidden, second forbidden, \\etc., types~\\cite{Beh82,Krm80}. %; see, for instance Refs.~\\cite{Ebe75,Krm80}.} There are exceptions, of course, as for instance, is the recent work of Marketin \\etal~\\cite{Mar09} on muon capture, where  the Walecka's classification  was used.  The formalism worked out by  Kuramoto~\\cite{Kur90}  is also frequently used for the evaluation of neutrino-nucleus cross-sections. It is   simpler than that of Walecka, but it does not contain   relativistic matrix elements, nor  is applicable for  muon capture rates.  More recently, we have introduced another formalism~\\cite{Krm02,Krm05,Sam06}. Besides of being   almost as simple as that of Kuramoto, it retains  relativistic terms and can be used for $\\mu$-capture. This formalism is briefly sketched here, including the consequences of the violation of the Conserved Vector Current (CVC) by the Coulomb field. It is furthermore simplified by  classifying the nuclear matrix elements in natural and unnatural parities. We also show how within the present formalism both the sum rule approach, and the formula for ERL  look like.  In Section \\ref{Sec2} we briefly describe  the formalism used to evaluate different  weak interacting processes. Some details are delegated to the Appendices: \\ref{A}) Contributions of natural and unnatural parity states to the transition rates, \\ref{B}) Sum rule approach for the inclusive neutrino-nucleus cross section, \\ref{C}) Formula for the inclusive neutrino-nucleus cross section at the extreme relativistic limit, and \\ref{D}) Formula for the muon capture rate. In Section \\ref{Sec3}  we present, and discuss the numerical results. Comparisons with experimental data, as well as with previous theoretical studies, are  done   whenever possible. Here we firstly sketch the two theoretical frameworks, namely the PQRPA and RQRPA,  used in the numerical calculations. In subsections \\ref{Sec3A}, and \\ref{Sec3B}  we present the results for the exclusive and inclusive processes, respectively. Finally, in Section \\ref{Sec4} we give a brief summary, and final conclusions. ",
        "conclusions": "\\label{Sec4}   The present work is a continuation of our previous works~\\cite{Krm02,Krm05}. In fact, the formalism for weak interaction processes introduced there is now elaborated more thoroughly yielding very  simple expressions for the transition rates,  which  greatly facilitate the numerical calculation. This is done through the separation of the nuclear matrix elements into their real and imaginary parts, which, in turn,  permits to split  the transition rates, for neutrino-nucleus reactions \\rf{2.22} into natural \\rf{A6} and unnatural parity  \\rf{A7} operators. Similar separation is done for  the  muon capture transition rate \\rf{D3} in Eqs.  \\rf{D4} and  \\rf{D5}. Moreover, consequences of explicit violation of CVC hypothesis by the Coulomb field \\rf{2.18} are addressed for the first time, and the sum rule approach for the inclusive cross section, proper  to the present formalism, has been worked out in the Appendix B. For the sake of completeness, the extreme relativistic limit of neutrino-nucleus cross section  is also presented in the Appendix C, where in the formula for transition rates turn out to be still simpler. We note that, except at very low neutrino energies, they can be used without any restriction in practical applications.  We have  discussed in details the inclusive properties that comprise:  i) Ground state energies in  $^{12}$B and $^{12}$N, and the corresponding GT $B$-values (Figure ~\\ref{F1}),  ii) Exclusive $^{12}$C($\\nu,e^-)^{12}$N cross-section $\\sigma_e(E_\\nu,1^+_{1})$, as a function of the incident neutrino energy $E_\\nu$ (Figures~\\ref{F2}, and ~\\ref{F3}),  iii) Exclusive $^{12}$C(${\\tilde{\\nu}},e^+)^{12}$B cross-section $\\sigma_{e^+}(E_{\\tilde{\\nu}},1^+_{1})$, as a function of the incident antineutrino energy $E_{\\tilde{\\nu}}$ (Figure ~\\ref{F4}), and  iv)  Muon capture transition rate to the $^{12}$B ground state,  and electron and muon folded cross-sections to the $^{12}$N ground state $\\overline{\\s}_\\e(1^+_1)$, and  $\\overline{\\s}_\\mu(1^+_1)$ (Figure ~\\ref{F5}).  \\noindent Special attention was paid to the interplay between the size of the configuration space, and the magnitude of the residual interaction within the $pp$-channel. It was found that as the first becomes larger, the second has to increase to obtain the agreement with the experimental data for the exclusive observables.   The main purpose of our discussion of exclusive properties was to put in evidence  the limitations of the RPA and the QRPA models. The basic problem in the implementation of  the  RPA is the lack of pairing correlations, i.e. the inability for opening the $1p_{3/2}$ shell, while deficiency of the standard QRPA is in the non-conservation of the number of particles, as evidenced by the wave functions \\rf{3.1}, \\rf{3.2}, and \\rf{3.3} presented in Sec.~\\ref{Sec3}. In this way  we have definitively established that the SM and  the PQRPA are proper theoretical frameworks to describe the ground state properties of $^{12}$B and  $^{12}$N. \\footnote{ After  our work has been  finished, Cheoun \\etal~\\cite{Ch10} have presented a new  evaluation of the ECS in $^{12}$C within the QRPA. They  get  good agreement with data for $\\overline{\\sigma}_e(^{12}$N), which is at variance with the previous QRPA calculation~\\cite{Vol00}.}  The inclusive  cross-sections $^{12}$C($\\nu,e^-)^{12}$N and $^{12}$C($\\tilde{\\nu},e^+)^{12}$B have been studied within the PQRPA in the same manner as the exclusive  ones for $E_{\\nu_e}$ up to $300$ MeV. As there are  no experimental available in this case the  comparison is done with the previous calculations  only \\footnote{ As already mentioned in the introduction the only available experimental data on $^{12}$C ICS  is the low neutrino energy  ($E_{\\nu_e}<60$ MeV) folded one, which has already been discussed in our previous works \\cite{Krm05,Sam06,Paa07}.}, and displayed  in Figures \\ref{F6}, and \\ref{F7}. Here, unlike within Figures \\ref{F2}, \\ref{F3}, and \\ref{F4}, we also show the results obtained with the other RPA-like models ~\\cite{Con72,Vol00,Kol99b,Paa07}, which could  be  a suitable framework for  describing   global nuclear properties such as it is the inclusive  cross-sections.  When the  size of the configuration space is enlarged the  calculated PQRPA cross-sections, at difference with the exclusive ones, steadily increase, and particulary for neutrino energies larger than $200$ MeV, in spite of including the particle-particle interaction.  At low energy they approach to  the cross section of the first-forbidden sum-rule limit, but are significantly smaller at high energies both for neutrino and antineutrino.  The largest space that we can deal within  the number projection procedure is the one that includes all the orbitals until the $N=6$ HO shell. This is the reason why we have recurred to the RQRPA where it is possible to employ larger configuration spaces. It seems that when the number of shells is increased to $N={30}$, and the  cut-off energy $E_{2qp}$ is large enough, the cross sections very likely converge as shown in Figures \\ref{F8}, \\ref{F9},  \\ref{F10}, and  \\ref{F11}.  The Figure \\ref{F10} also indicates  that  the RQRPA is a promising nuclear model to reproduce the quasi elastic ($\\nu_\\mu,^{12}$C) cross section in the region of $E_{\\nu_\\mu}\\sim 1$ GeV which has been measured recently at MiniBooNE~\\cite{MiniBooNE}. We do not know whether the discrepancy between the experiment and the theory comes from the non completeness of the configuration space or from the smallness of the effective axial-vector coupling constant that we are using $\\gA=1$. It could also happens that we need $\\gA=1$ for the low energy exclusive cross section and $\\gA=1.255$ for the high energy inclusive cross section. We do not understand the reason for such a energy dependence of $\\gA$, but it is consistent with the Eq. (23) in Ref.~\\cite{Na82} where it is shown that  for the low energy $\\beta$-decay $\\gA$ could be much more quenched that the total GT strength. We hope to be able to say more on this matter in the next future.  We have also addressed the issue of multipole composition of the inclusive cross sections, by separating them into allowed ($J^\\pi=0^+,1^+$), first-forbidden ($J^\\pi=0^-,1^-,2^-$), second-forbidden ($J^\\pi=2^+,3^+$), and third-forbidden ($J^\\pi=3^-,4^-$) processes. The results for  the antineutrino reaction $^{12}$C($\\tilde{\\nu},e^+)^{12}$B are displayed in Figure~\\ref{F12} both for the PQRPA and the RQRPA. Of course, similar  results are obtained also for neutrinos. We remark  that the spectral functions $d\\sigma^A_{e^+}(E_{\\tilde{\\nu}})/dE_{\\tilde{\\nu}}$, when evaluated within the PQRPA, clearly put into evidence the resonant structure of the allowed cross-section, which is mainly of the GT type.  The study of the partial ICSs has been related with the proposal done in Refs.~\\cite{Laz07,Vol04} of performing nuclear structure studies of forbidden processes by using low energy neutrino and antineutrino beams. From the results shown in Table \\ref{table} for the flux-averaged cross sections $\\overline{\\s}_{e^+}$ in the reaction $^{12}{\\rm C}({\\tilde{\\nu}},e^+)^{12}{\\rm B}$ we show  that the contribution of allowed transitions decreases gradually in favor of the first forbidden transitions according with the increase of $\\gamma$-boost. %{\\color{blue} We conclude  that to study high forbidden reactions one would need  ${\\tilde{\\nu}}$-fluxes with $E_{\\tilde{\\nu}}$ up to  $\\gtrsim 150$ MeV in $^{12}C$. %}   At the end we considered possible astrophysical applications of the $\\nu/\\tilde{\\nu}$-$^{12}$C nucleus folded  cross sections $\\overline{\\s}_{e^\\pm}$, using supernovae $\\nu/\\tilde{\\nu}$ spectra represented by a normalized Fermi-Dirac distribution with mean energy $\\langle E_\\nu \\rangle \\approx 3.15 \\times T_\\nu$, and zero chemical potential.  It is found that for temperature $T_\\nu=3-5$ MeV both the PQRPA and  RQRPA models yield significantly larger  cross sections that the previously used shell model."
    },
    "1005/1005.4903_arXiv.txt": {
        "abstract": "We describe a survey in the ELAIS N2 region with the VLA at 43.4~GHz, carried out with 1627 independent snapshot observations in D-configuration and covering about 0.5 \\sdeg. One certain source is detected, a previously-catalogued flat-spectrum QSO at z=2.2. A few ($<5$) other sources may be present at about the $3\\sigma$ detection level, as determined from positions of source-like deflections coinciding with blue stellar objects, or with sources from lower-frequency surveys.  Independently we show how all the source-like detections identified in the data can be used with a maximum-likelihood technique to constrain the 43-GHz source counts at a level of $\\sim$7~mJy. Previous estimates of the counts at 43 GHz, based on lower-frequency counts and spectral measurements, are consistent with these constraints, although the present results are suggestive of somewhat higher surface densities at the 7~mJy level. They do not provide direct evidence of intrusion of a previously unknown source population, although the several candidate sources need examination before such a population can be ruled out. ",
        "introduction": "We present a survey to detect extragalactic radio sources at 43 GHz with repeated independent pointings of the D-configuration VLA in snapshot mode.  The survey is centered on ELAIS-N2 (RA~16h~36m~48s, Dec +41\\deg~01'~45'', J2000), and covers about 0.5 \\sdeg~down to 7~mJy. The primary objective was to establish the 43-GHz source surface density at mJy levels, an issue of importance in assessing the discrete-source foreground contamination of the CMB signal.  It is only recently that we have gained insight into what populates the radio sky at survey frequencies above 5~GHz. By this frequency some 60 per cent of the sources at Jy levels have `flat' spectra ($\\alpha > -0.5$, where flux density $S$ is related to frequency $\\nu$ via $S \\propto \\nu^\\alpha$) over some region of the radio domain. This is the signature of optically-thick synchrotron emission and compact structure, so-called `flat-spectrum' sources. (Surveys at frequencies above 150~GHz do find inverted-spectrum sources presumably dominated by free-free emission, probably dust emitters related to or part of the sub-mm galaxy population. These do not appear at frequencies as low as 43~GHz -- see Viera et al. 2009). `Flat-spectrum' objects do not generally have $\\alpha = 0$ over a large frequency range; rather their spectra show inversions, points of inflection, and sometimes single peaks (e.g. Gigahertz-Peaked Spectra or GPS sources). The optical counterparts of the great majority of these sources are stellar in appearance and are Flat-Spectrum Radio QSOs (FSRQs) or BL\\,Lac objects.  How rapidly does this trend to such objects and away from conventional steep-spectrum -- generally optically-thin spectrum radio galaxies -- continue with increasing frequency above 5~GHz? When and at what flux-density level do we find new populations of flat or inverted-spectrum objects?  The pioneering survey of Brandie \\& Bridle (1974) covered 0.8~sterad at 8~GHz and found 55 sources, of which ``more than 2/3'' were stellar objects, i.e. QSOs or BL\\,Lac objects. Subsequent deep surveys by Windhorst et al. (1993) enabled definition of an 8.4-GHz source count down to 10~$\\mu$Jy.  Brandie and Bridle found that at their relatively high flux densities, 70 per~cent of the sources have flat or inverted spectra. Some spectrally-extreme sources certainly exist -- Edge et al. (1998) found two with spectral peaks at $\\ge$~20~GHz -- but the proportion of such objects to be detected in very high frequency surveys remains unknown.  The first ventures at higher frequencies were to search for contaminants to CMB experiments, with the Ryle Telescope  survey at 15~GHz leading the way. Taylor et al. (2001) found 66 sources in 63 \\sdeg~above 20~mJy. The continuation of this came to be the Cambridge 9C survey (Waldram et al. 2003; Bolton et al. 2004; Waldram et al. 2009). Some 465 sources were catalogued above 25 mJy in 520 \\sdeg. Examination of a subset of these sources (Bolton et al. 2004) revealed that (a) at a flux limit of 25~mJy 18 per cent show spectra that peak at 5 GHz or above; and (b) at 60~mJy 27 per cent peak at 5~GHz or above, and increasing the flux limit increases the fraction of rising-spectrum sources.  (The statistical dependence of radio spectrum on flux-density level has been discovered relatively regularly since 1970; see e.g. Kellermann \\& Wall 1987.) Colours and the stellar nature of optical counterparts indicate that between a third and a half of the peaked-spectrum sources are QSOs.  The major contribution to surveys and samples at these higher frequencies has come from AT20G, the large-area 20-GHz survey with the Australia Telescope Compact Array (Ricci et al 2004; Sadler et al. 2006; Massardi et al. 2008). By 2007 the survey had covered $-90$\\deg\\ $ < \\delta < -15$\\deg\\ cataloguing 4400 sources down to $\\sim$50~mJy; another 1500 are expected in the final survey zone $-15$\\deg$ < \\delta < 0$\\deg. These catalogues will render incorrect the statements made at the start of each of these papers about our poor knowledge of the high-frequency radio sky. Massardi et al. studied the 320 brightest sources from the 2007 version of the catalogue. The major conclusions do not differ dramatically from those from 9C: (a) most sources do not show power-law spectra; (b) the spectral complexity ensures that it is impossible to select a low-frequency sample which will constitute a complete high-frequency sample; (c) spectral steepening is common and correlates with redshift, perhaps due to changing rest frame frequency; (d) 77 per cent of the brightest-source sample are QSOs, 19 per cent are galaxies, and the remainder are blank on the SuperCosmos Sky Survey (SSS; see {\\it www-wfau.roe.ac.uk/sss/}) scanned versions of the UKSTU plates. The galaxies have a median $B_J$ = 17.7, the QSOs 18.6 -- these medians lie far above the plate limits.  Further characterization of the source population at high frequencies has come from multi-frequency flux-density measurements of sources detected in the WMAP (Wilkinson Microwave Anisotropy Probe) CMB surveys. WMAP has produced complete samples of the brightest radio sources over the sky to a flux-density limit of about 2 Jy at each of its primary frequencies 23, 33, 43, 61 and 94 GHz (Hinshaw et al. 2007; L\\'{o}pez-Caniego et al. 2007; Wright et al. 2009). With measurements at 16 and 33 GHz using the VSA and AMI, Davies et al. (2009) and Franzen et al. (2009) found the following for a sample complete to 1.1~Jy at 33 GHz: (1) the proportion of flat-spectrum objects is very high: 93 per cent have spectra with $\\alpha^{34}_{14~\\rm{GHz}} > -0.5$; (2) 44 per cent have $\\alpha^{34}_{14~\\rm{GHz}} > 0.0$, i.e. rising spectra; (3) the variable-source bias (a selection effect which mimics an Eddington bias; see Wall et al. 2005; Wall 2007) is in clear evidence from comparing repeated WMAP and AMI 33-GHz flux densities; (3) on timescales of $\\sim$1.5yr, 20 per cent of sources varied by more than 20 per cent at 33 GHz with a high degree of correlation between variation at 16 and 33 GHz; and (4) variable sources have a marginally flatter spectrum ($\\overline{\\alpha} = -0.06\\pm0.05$) than non-variable sources ($\\overline{\\alpha} = -0.13\\pm0.04$).   Sadler et al. (2008) attempted to characterize the populations to be found at yet higher frequencies. With flux densities measured for 130 sources at 95 GHz and 20~GHz with the ATCA, the resultant distribution of spectral indices together with source counts from AT20G was used to predict the source count at 95 GHz above a level of 100 mJy. This prediction is lower than that from models of luminosity functions and evolution (De Zotti et al. 2005). The authors conclude that this is due to the fact that the majority of sources with flat or rising spectra over the range 5 to 20~GHz show a spectral turnover between 20 and 95~GHz.  It is the purpose of this paper to describe a limited survey at mJy levels at 43.4 GHz with the VLA, and to use this to construct a source count at this frequency. We want to see if there is evidence for emergent new populations that might affect future CMB experiments such as Planck. In addition we wish to track the space density of the different classes of object which emerge as survey frequency is raised, and for this, a source count is an essential complement to radio and optical spectral data. Descriptions of this space density for {\\em combined} populations of flat-spectrum objects have been obtained (e.g De Zotti et al. 2005; Wall et al. 2005; Ricci et al. 2006), but not for the individual (e.g. GPS) populations.  With regard to a source count at 43 GHz, the WMAP all-sky survey at 43~GHz defined the high-flux-density end of the count as well as it can ever be done. To extend the 43-GHz count to lower flux levels, Waldram et al. (2007) used their 15-GHz 9C survey plus spectral information to predict source counts at this frequency (as well as at several other frequencies). At a somewhat lower frequency, Mason et al. (2009) measured flux densities at 31 GHz for 3165 NVSS sources using the Green Bank 100m and the Owens Valley 40m telescopes. From these measurements they projected an integral source count at 31~GHz in the range 1 to 4~mJy, and found that the surface densities implied account for $21 \\pm 7$ per cent of the amplitude of the power detected in excess of intrinsic anisotropy by the Cosmic Background Imager (CBI) at $\\ell > 2000$. We return to these results in \\S~\\ref{scest} and \\S~\\ref{disc}.  \\begin{figure*} \\vspace{9cm}  \\special{psfile=kband_grid_EN2_FIRST.eps hoffset=45 voffset=245 vscale=70 hscale=70 angle=-90}  \\caption{The observing grid (Dec vs. RA, J2000) for our 1627 independent snapshot observations with the VLA at 43 GHz. The green circles represent the approximate size of the synthesized fields; the field-centre separation is twice the FWHM of the primary beam. Red dots mark source positions from the FIRST catalogue (Becker et al. 1995); black dots show source positions from the ELAIS N2 1.4-GHz catalogue (Ciliegi et al. 1999).}  \\label{obsgrid}  \\end{figure*}       Surveying at frequencies as high as 43~GHz is not straightforward, as this pilot study will demonstrate. Hence this paper does not follow the  conventional survey, analysis, catalogue, source-count etc. route --  the data do not lend themselves to such a procedure. We thus present a road map for orientation. \\begin{enumerate} \\item \\S~\\ref{surdes} {\\em Survey Design} considers why we chose the region ELAIS N2 and how we came to the plan for coverage in terms of independent snapshots of given cycle parameters. The Appendix A1 describes an analysis leading to this design. \\item {\\em Observations and Reductions} (\\S~\\ref{obsred}) constitutes the conventional part of the paper. We used standard AIPS analysis routines. \\item The 1627 independent snapshot images show many source-like deflections, as decribed in \\S~\\ref{obsred}. With our lack of knowledge of the sky at 43~GHz, we try in \\S~\\ref{dets}, {\\em Finding Real Sources} to use as many avenues as possible to explore the reality of these deflections. To do so, we select a sample of deflections from all snapshot images down to a level  of 3 times the rms noise level on each image. It is essential to emphasize at the outset that {\\em we never use a 3$\\sigma$ sample in a source-counting process}, as the effects of Eddington bias are lethal at such levels. However the sample gives us the broadest scope to establish the frequency of real detections, and we try to do this by examining the statistics of the sample, the coincidences of deflections with known sources from radio catalogues at other frequencies, and the coincidences with optical counterparts using the Palomar II Sky Survey and the Sloan Digital Sky Survey. \\item The minimal returns from the exercises of \\S~\\ref{dets} led us to consider how to use the survey data to estimate the source count at 43~GHz. We do so in \\S~\\ref{scest} by analyzing the imprint of the primary beam shape on the expected radial distribution of sources in the collective images, and we develop a technique for this that uses {\\em each deflection}, real or otherwise, in a maximum-likelihood analysis. \\item Finally (\\S~\\ref{disc}) we compare the results of this analysis with estimates of counts at other mm-wavelength frequencies. \\end{enumerate} ",
        "conclusions": "\\label{disc} 1. This 43-GHz survey (VLA D-array) reaches approximately 7 mJy over an area of 0.5 \\sdeg, with a primary beam of 1.0~arcmin FWHM and a circular synthesized beam of 2.3~arcsec FWHM. Done in snapshot mode to follow the `wide-shallow' dictum, it found many apparent sources, most of which on inspection proved to be at the junction of response lines. Most can be eliminated on this basis, and on the basis that they are extended to the 2.3 arcsec synthesized beam. The radial distribution likewise makes it evident that the overwhelming majority of the deflections are not real sources.  2. However the survey does detect real sources, the flat-spectrum QSO 16382+414 for certain, and some additional not-very-probable candidates (Table~\\ref{coinc_all}) for which further observations are required. One of these is coincident in position with previously-catalogued sources at 1.4~GHz; there is no candidate identification above the limits of the SDSS. Of the 6 other `sources' in the table, all are coincident with blue stellar objects close to the region occupied by $z \\sim 2.5$ FSRQ. These putative 6 sources all must have extreme inverted spectra because stacking suggests that the average flux density at 1.4~GHz must be below 0.1 mJy.  3. We have developed a likelihood method to estimate source counts in a synthesis survey which has thrown up a combination of real and random sources. The method examines the radial distribution in each synthesized field, and because the real sources should decrease with radial distance from the field centre in a predictable manner, in principle both the source-count normalization and the source count power-law index can be determined. With this process, neither detailed cataloguing of the sources nor {\\it a priori} decisions about which deflections are real are necessary. We have applied this method to the present data, but the dominance of random deflections (by a ratio of perhaps 10000:1) is such that only an upper limit to the surface density of real sources can be derived. No information on the slope of the count is available from the present data; we have adopted a slope estimate for the appropriate flux-density range from Waldram et al. (2007), although the resulting upper-limit contours are relatively insensitive to this.  4. The present results are consistent with extrapolating the count estimate of Waldram et al. (2007), and with the 31-GHz count estimate of Mason et al. (2009), but suggestive of a somewhat higher source surface density than found in either investigation. Such a surface density enhancement might account for the excess signal at $\\ell > 2000$ found by CBI. We note that the De Zotti et al. (2005) count models, calculated at high frequencies from detailed synthesis of populations from lower-frequency surveys, lie significantly above the count estimates of Waldram et al. (2007).  5. In terms of individual sources, the 43-GHz survey yield is small and provides no {\\it direct} evidence that any new types of inverted-spectrum sources appear at mJy levels. The lower bound of our error box in Fig.~\\ref{sc43} is strongly dependent on the certain detection of a single source; if any of our candidate sources of Table~\\ref{coinc_all} should prove to be real, then a count enhanced above the Waldram et al. (2007) estimate becomes probable. Any such sources will have extremely inverted spectra, as we and Mason et al. (2009) have noted.  6. The observations indicate that there are considerable challenges involved in a snapshot synthesis survey at a frequency as high as 43 GHz.   \\vspace{0.5cm}  \\noindent {\\bf Acknowledgements} We thank Chris Willott for helpful discussions about field choice. We are very grateful to Barry Clark for his efforts in scheduling the VLA to accommodate a late improvement in our observing proposal. Jim Condon proferred useful advice. Eric Greisen gave much help and advice with running AIPS and checking systems remotely. Melanie Gendre helped with running some of the AIPS procedures. Ron Ekers offered many constructive comments on an early draft. Chris Blake gave useful advice on SDSS data. We thank the referee for helpful comments. JVW would particularly like to thank Katherine Blundell for help with computational issues and and with running AIPS over the past several years. JVW acknowledges support during the course of this work through Canada NSERC Discovery Grants.  The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.  This research has made use of data obtained from the SuperCOSMOS Sky Surveys (SSS), prepared and hosted by the Wide Field Astronomy Unit, Institute for Astronomy, University of Edinburgh, which is funded by the UK Science and Technology Facilities Council. The SSS Web Site is {\\it www-wfau.roe.ac.uk/sss/}.  The research also made use of data from the SDSS survey. Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS Web Site is {\\it www.sdss.org/}.  \\appendix"
    },
    "1005/1005.3955.txt": {
        "abstract": "\\vskip - 2.2in \\rightline{UMN--TH--2832/10} \\rightline{FTPI--MINN--10/03} \\rightline{January 2010} \\vskip 1.4in Four lectures on Big Bang cosmology, including microwave background radiation, Big Bang nucleosynthesis, dark matter, inflation, and baryogenesis. ",
        "introduction": "\\subsection{The FLRW metric and its consequences}  The standard Big Bang model assumes homogeneity and isotropy.  As a result,  one can construct the space-time metric by embedding a maximally symmetric three dimensional space in a four dimensional space-time (see, e.g., Ref.~\\cite{wein}). The most general form for a metric of this type is the Friedmann--Lema\\^{i}tre--Robertson--Walker metric which in co-moving coordinates is given by \\beq ds^2  = dt^2  -  R^2(t)\\left[ {dr^2 \\over \\left(1-kr^2\\right) } + r^2 \\left(d\\theta^2  + \\sin^2 \\theta d\\phi^2 \\right)\\right]\t, \\label{met} \\eeq where $R(t)$ is the cosmological scale factor and $k$ the three-space curvature constant ($k = 0, +1, -1$ for a spatially flat, closed or open universe). $k$ and $R$ are the only two quantities in the metric which distinguish it from flat Minkowski space. It is  also common to assume the perfect fluid form for the energy-momentum tensor \\beq T^{\\mu\\nu}   = pg^{\\mu\\nu}   + (p + \\rho)u^\\mu u^\\nu \t\t, \\eeq where $g_{\\mu\\nu}$   is the space-time metric described by Eq.~(\\ref{met}), $p$ is the isotropic pressure, $\\rho$ is the energy density and $u^\\mu  = (1,0,0,0)$ is the velocity vector for the isotropic fluid.  The $(00)$ component of Einstein's equation \\beq R_{\\mu\\nu} - \\frac{1}{2}g_{\\mu\\nu}R - \\Lambda g_{\\mu\\nu} = 8 \\pi G_N T_{\\mu\\nu} \\eeq yields the Friedmann equation \\beq H^2  \\equiv \\left({\\dot{R} \\over R}\\right)^2  = { 8 \\pi G_N \\rho \\over 3} - { k \\over R^2}  + {\\Lambda \\over 3} , \\label{H} \\eeq and the $(ii)$ components give \\beq \\left({\\ddot{R} \\over R}\\right) = {\\Lambda \\over 3} - { 4 \\pi G_N ( \\rho + 3p) \\over 3} , \\label{ii} \\eeq where $\\Lambda$ is the cosmological constant. In addition, from ${T^{\\mu\\nu}}_{;\\nu}   =  0$, we obtain \\beq \\dot{\\rho} = -3H(\\rho + p)\t. \\label{rhod} \\eeq Note that of these last three equations, only two are actually independent. These equations form the basis of the standard Big Bang model.  At early times ($t < 10^5 $ yrs) the Universe is thought to have been dominated by radiation so that the equation of state can be given by $p = \\rho/3$.  If we neglect the contributions to $H$ from $k$ and $\\Lambda$ (this is always a good approximation for small enough $R$) then we find that \\beq R(t) \\sim t^{1/2} \\qquad \\eeq and $\\rho \\sim R^{-4}$  so that $t \\sim (3/32 \\pi G_N\\rho)^{1/2}$. Similarly for a matter or dust dominated universe with $p = 0$, \\beq R(t) \\sim t^{2/3} \\eeq and $\\rho \\sim R^{-3}$.  The Universe makes the transition between radiation and matter domination when $\\rho_{rad} = \\rho_{matter}$ or when $T \\simeq$ few $\\times~10^3$ K. In a vacuum or $\\Lambda$ dominated universe (that we are approaching today) \\beq R(t) \\sim e^{\\sqrt{\\Lambda/3}\\  t} . \\label{Vdom} \\eeq   More general solutions for the behaviour of the scale factor can easily be found by defining a quantity $Q$: \\beq Q = \\frac{3k}{R^2} - 8\\pi G_N \\rho . \\eeq If we further assume $p = (\\gamma-1) \\rho$ (with $\\gamma$ between 1 and 2), we have that $\\rho \\sim R^{-3\\gamma}$ and we can write \\beq \\frac{\\dot R}{R} = \\pm \\left[ \\frac{\\Lambda - Q}{3} \\right]^{1/2} , \\label{Qevol} \\eeq implying that $Q \\le \\Lambda$. For all choices of $k$, the function $Q \\to -\\infty$ as $R \\to 0$. When $k = +1$, it is easy to see that $Q$  has  a maximum (with $Q > 0$) and then tends to 0 as $R \\to \\infty$. When $k = 0$ or $-1$, $Q$ monotonically increases towards 0 as $R \\to \\infty$. In Fig.~\\ref{fig:Q}, the function $Q$ is plotted qualitatively as a function of the scale factor.   \\begin{figure} \\centering\\includegraphics[width=.5\\linewidth]{Qs} \\caption{The qualitative behaviour of the scale factor $R$ as obtained using the condition $Q \\le \\Lambda$ for the three choices of the curvature constant $k$} \\label{fig:Q} \\end{figure}   Let us first consider the more interesting case of $k = +1$. For a cosmological constant $\\Lambda  > Q_{max}$, $Q< \\Lambda$ for all $R$. In this case, two distinct solutions are possible: the universe may expand forever from a singularity at $R = 0$ to  $R = \\infty$, or by choosing the lower sign in Eq.~(\\ref{Qevol}), the universe may collapse from infinity to a singularity.  There are also two possible solutions for $0 < \\Lambda < Q_{max}$, the case depicted in Fig.~\\ref{fig:Q}. The universe may again start at a singularity at $R=0$ and expand to the point where $Q = \\Lambda$ and then recollapse.  Alternatively, the Universe may start from infinity and collapse to the point where $Q =  \\Lambda$ (at larger $R$),   bounce back and expand to infinity. When $\\Lambda = Q_{max}$ there are a total of five solutions. At the value of $R$ such that $Q = Q_{max}$, we have the Einstein Static Universe. This is the only solution which is neither expanding nor contracting.  The remaining four solutions either asymptotically expand or contract towards or away from the Einstein Static case. Finally for $\\Lambda < 0$, there is only one solution for which the universe expands and subsequently recollapses.  Finding the solutions when $k = 0$ or $-1$ is similar, and the behaviour of $R$ depends only on whether $\\Lambda$ is positive or negative. For $\\Lambda \\ge 0$ there are two solutions, for $\\Lambda < 0$, there is only one. When $\\Lambda = 0$, we obtain the standard notion that closed universes recollapse, while open and flat universes expand forever. When $\\Lambda \\ne 0$, these simple associations are spoiled as a closed universe can expand forever (for large enough $\\Lambda$) and open and flat universes can recollapse (for $\\Lambda < 0$).  Exact solutions to the equations of motion can be obtained relatively easily in terms of conformal coordinates. We can, for example, rewrite the metric as \\beq ds^2  = dt^2  - R^2(t)\\left[ {d\\chi^2 } + f^2(\\chi) \\left(d\\theta^2  + \\sin^2 \\theta d\\phi^2 \\right)\\right]\t , \\eeq where \\beq f(\\chi) = \\left\\{ \\begin{array}{ll} \\sinh \\chi & k = -1\\\\ \\chi & k = 0 \\\\ \\sin \\chi & k = +1 \\end{array} \\right. . \\eeq We can go further and define a conformal time coordinate using $R d\\eta = dt$ so that \\beq ds^2  = R^2(\\eta )\\left[  dt^2  -  {d\\chi^2 } + f^2(\\chi) \\left(d\\theta^2  + \\sin^2 \\theta d\\phi^2 \\right)\\right]\t . \\eeq In terms of these coordinates, the Friedmann equation becomes \\beq R^{\\prime\\prime} + kR =\\frac{4 \\pi G_N}{3} (\\rho - 3p) R^3 , \\eeq where $^\\prime$ denotes a derivative with respect to $\\eta$ and is easily solved \\beq R \\propto \\left\\{ \\begin{array}{ll} \\cosh \\eta -1 & k = -1\\\\ \\eta^2/2 & k = 0 \\\\ 1 - \\cos \\eta & k = +1 \\end{array} \\right.  \\qquad t \\propto  \\left\\{ \\begin{array}{ll} \\sinh \\eta -\\eta & k = -1\\\\ \\eta^3/6 & k = 0 \\\\ \\eta - \\sin \\eta & k = +1 \\end{array} \\right. \\eeq for $p=0$ and \\beq R \\propto \\left\\{ \\begin{array}{ll} \\sinh \\eta  & k = -1\\\\ \\eta & k = 0 \\\\ \\sin \\eta & k = +1 \\end{array} \\right.  \\qquad t \\propto  \\left\\{ \\begin{array}{ll} \\cosh \\eta -1 & k = -1\\\\ \\eta^2/2 & k = 0 \\\\ 1 - \\cos \\eta & k = +1 \\end{array} \\right. \\eeq for $p= \\rho/3$.  In the absence of a cosmological constant, one can define a critical energy density $\\rho_c$ such that $\\rho =\\rho_c$  for $k = 0$ \\beq \\rho_c  = 3H^2 / 8 \\pi G_N\t\t. \\eeq In terms of the present value of the Hubble parameter this is \\beq \\rho_c  = 1.88 \\times 10^{-29} {h_0}^2  {\\rm ~g~cm}^{-3}  , \\eeq where \\beq h_0  = H_0 /(100 {\\rm ~km~Mpc}^{-1}   {\\rm ~s}^{-1}  )\t\t. \\eeq The cosmological density parameter is then defined by \\beq \\Omega \\equiv {\\rho \\over \\rho_c} \t. \\eeq It is useful to also define a deceleration parameter \\beq q_0 = -\\frac{{\\ddot R}_0 R_0}{{\\dot R}^2_0} . \\eeq This standard definition was formulated under the presumption that the expansion rate of the Universe is in fact slowing down.  As noted above, and discussed further below, modern measurements indicate the opposite. That is, the expansion is accelerating (meaning that $q_0 < 0$). The $(ii)$ component, Eq.~(\\ref{ii}), can be written in terms of $q_0$ as \\beq -2 q_0 H_0^2 = \\frac{2 \\Lambda}{3} - \\frac{8 \\pi G_N \\rho_0}{3} \\eeq when the pressure is neglected. This can be combined with the Friedmann equation, Eq.~(\\ref{H}), and rewritten as \\beq {k \\over R_0^2} = \\Lambda + H_0^2 (2q_0 - 1)\t , \\label{o-1} \\eeq or \\beq {k \\over R_0^2} =  H_0^2 (\\frac{3}{2}\\Omega_0 -q_0 - 1)\t. \\eeq Furthermore, when $\\Lambda = 0$, $q_0 = \\Omega_0/2$ so that $k = 0, +1, -1$ corresponds to $\\Omega = 1, \\Omega > 1$ and $\\Omega < 1$. Observational limits on $h_0$ and $\\Omega$ are~\\cite{wmap} \\beq h_0 = 0.71 \\pm 0.01  \\qquad \\Omega_0 = 1.006 \\pm 0.006. \\label{range} \\eeq  It is important to note that $\\Omega$ is a function of time or of the scale factor. The qualitative evolution of $\\Omega$ is shown in Fig.~\\ref{fig:myfig} for $\\Lambda = 0$. For a spatially flat Universe, $\\Omega = 1$ always. When $k = +1$, there is a maximum value for the scale factor $R$. At early times (small values of $R$), $\\Omega$ always tends to one. Note that the fact that we do not yet know the sign of $k$, or equivalently whether $\\Omega$ is larger than or smaller than unity, implies that we are at present still at the very left in the figure.  What makes this peculiar is that one would normally expect the sign of $k$ to become apparent after a Planck time of $10^{-44}$ s. It is extremely puzzling that some $10^{60}$ Planck times later, we still do not know the sign of $k$.   \\begin{figure} \\centering\\includegraphics[width=.7\\linewidth]{omega3s} \\caption{The evolution of the cosmological density parameter $\\Omega$ as a function of the scale factor for a closed, open and spatially flat Universe} \\label{fig:myfig} \\end{figure}  The Friedmann equation also lends itself to integration to determine the age of the Universe.  Note that for a given equation of state, we can write $\\rho = \\rho_0 (R_0/R)^{3\\gamma}$, and \\beq H_0 t = \\int_0^1 \\frac{dx}{\\left[ 1 - \\Omega_0 - \\Omega_\\Lambda + \\Omega_0 x^{2 - 3\\gamma}+ \\Omega_\\Lambda x^2  \\right]^{1/2} } , \\eeq where $\\Omega _\\Lambda = \\Lambda/3 H_0^2$. When $\\Lambda = 0$ and $\\Omega_0 = 1$, this is easily integrated to give \\beq t = \\frac{2}{3 H_0}  \\qquad \\gamma = 1 \\eeq or \\beq t = \\frac{1}{2 H_0}  \\qquad \\gamma = \\frac{4}{3} . \\eeq  Because of the finite age of the Universe in the Big Bang model, there is a particle horizon corresponding to the maximum distance traversed by light. In general, we can write the proper distances between a location specified by comoving coordinate $t_1, r_1$ and our location at $t$ and $r=0$ in terms of the metric~(\\ref{met}) \\beq d_p = R(t) \\int_0^{r_1} \\frac{dr^\\prime}{\\sqrt{1 - k {r^\\prime}^2}} . \\eeq For light paths ($ds^2 = 0$) \\beq \\int_{t_1}^t \\frac{dt^\\prime}{R(t^\\prime)} = \\int_0^{r_1} \\frac{dr^\\prime}{\\sqrt{1 - k {r^\\prime}^2}} . \\eeq As $t_1 \\to 0$, $r_1$ becomes the maximum coordinate distance from which we can receive a signal. Thus the particle horizon is defined by \\beq d_H = R(t) \\int_0^{r_H} \\frac{dr^\\prime}{\\sqrt{1 - k {r^\\prime}^2}} = R(t) \\int_0^t \\frac{dt^\\prime}{R(t^\\prime)} . \\label{dh} \\eeq For $\\Lambda = 0$ and $\\Omega_0 = 1$, we again obtain very simple solutions: \\beq d_H  =3 t  \\qquad \\gamma = 1 \\eeq or \\beq d_H =2t  \\qquad \\gamma = \\frac{4}{3} . \\eeq Note that because $d_H/R$ grows with time, we see more of the Universe as time goes on. That is, new information is continuously entering our particle horizon.       \\subsection{The hot thermal Universe}  The epoch of recombination occurs when electrons and protons form neutral hydrogen  through $e^{ -  } +  p     \\rightarrow $  H  $+   \\gamma $ at a temperature $T_{ R} {~ \\sim }$~few $\\times 10^{3}$ K  ${~\\sim}~1$ eV.  For $T < T_{R}$, photons are decoupled while for $T > T_{ R}$,  photons are in thermal equilibrium and at higher temperatures, the Universe is radiation dominated and the content of the radiation plays a very important role.  Today, the content of the microwave background consists of photons with $T_0 =  2.725  \\pm 0.001$ K~\\cite{cobem}. We can calculate the energy density of photons from \\beq \\rho_\\gamma  = \\int E_\\gamma dn_\\gamma , \\label{rhog} \\eeq where the density of states is given by \\beq dn_\\gamma  =   {g_\\gamma \\over 2 \\pi^{ 2}}[exp(E_\\gamma/T)-1]^{ -  1} q^{ 2} dq , \\eeq and $g_\\gamma = $  2 simply counts the number of degrees of freedom for photons, $E_\\gamma =  q$ is just the photon energy (momentum). (I am using units such that $\\hbar =  c  = k_{ B}   =$  1 and will do so through the remainder of these lectures.) Integrating Eq.~(\\ref{rhog}) gives \\beq \\rho_\\gamma = {\\pi^2 \\over 15} T^4 , \\eeq which is the familiar blackbody result. In addition, we also have \\beq p = \\frac{1}{3} \\rho \\qquad s = \\frac{4}{3} \\frac{\\rho}{T} \\qquad n_\\gamma = \\frac{2 \\zeta(3)}{\\pi^2} T^3 . \\eeq  In general, at very early times, at very high temperatures, other particle degrees of freedom join the radiation background when $T{~\\sim }~m_{i}$  for each  particle type $i$ if that type is brought into thermal equilibrium through interactions.  In equilibrium, the energy density of a particle type $i$ is given by \\beq \\rho_{i}  = \\int E_{i} dn_{q_{i}} \\eeq and \\beq dn_{q_{i}} = {g_{i} \\over 2  \\pi^{ 2}}[\\exp[(E_{q_{i}} - \\mu_{i})/T] \\pm 1]^{ -1 }q^{2}dq , \\eeq where again $g_{i}$ counts the total number of degrees of freedom for type $i$, \\beq E_{q_{i}} =  \\left(m_{i}^{2} + q_{i}^{ 2}\\right)^{1/2} , \\eeq $\\mu_{i}$ is the chemical potential if present and  $ \\pm$ corresponds to either Fermi or Bose statistics.  In the limit that  $T \\gg m_{i} $  the total energy density can be conveniently expressed by \\beq \\rho {} = \\left( \\sum_B g_{B} + {7 \\over 8} \\sum_F  g_{F} \\right) {\\pi^{ 2} \\over 30}  T^{4}     \\equiv    {\\pi^{ 2} \\over 30} N(T) T^{4} , \\label{NT} \\eeq where $g_{B(F)} $  are the total number of boson (fermion) degrees of freedom and the sum runs over all boson (fermion) states with $m \\ll T$.  The factor of 7/8 is due to the difference between the Fermi and Bose integrals.  Equation~(\\ref{NT}) defines $N(T)$ by taking into account  new particle degrees of freedom as the temperature is raised.  In the radiation dominated epoch, Eq.~(\\ref{rhod}) can be integrated (neglecting the $T$-dependence of $N$) giving us a relationship between the age of the Universe and its temperature \\beq t = \\left({90 \\over 32 \\pi^3 G_{ N} N(T)}\\right)^{ 1/2}  T^{ -  2} . \\eeq Put into a more convenient form \\beq tT_{\\rm MeV}^{ 2}  = 2.4 [N(T)]^{-1/2}     . \\label{tT2} \\eeq where $t$ is measured in seconds and $T_{\\rm MeV} $  in units of MeV.  The value of $N(T)$ at any given temperature depends on the particle physics model.  In the standard $SU(3) \\times SU(2)\\times U(1)$  model, we can specify $N(T)$ up to temperatures of order 100~GeV. The value of $N$ in the Standard Model can be seen in Table~\\ref{tab:NT}.   \\begin{table} \\caption{Effective numbers of degrees of freedom in the Standard Model} %\\vspace{2pt} \\begin{center} \\begin{tabular}{llc} \\hline {\\bf Temperature} & {\\bf New particles} \\qquad &\\boldmath$4N(T)$ \\\\ \\hline\\rule{0pt}{12pt} $T < m_{ e}$   &     $\\gamma$'s +   $\\nu$'s & 29 \\\\ $m_{ e} <   T  < m_\\mu$ &    $e^{\\pm}$ & 43 \\\\ $m_\\mu <  T  < m_\\pi$  &   $\\mu {}^{\\pm}$ & 57 \\\\ $m_\\pi <  T < T_{ c}^{*}$  & $\\pi$'s & 69 \\\\ $T_{ c} <  T  < m_{\\rm charm}$ \\qquad & -  $\\pi$'s + $  u,{\\bar u},d,{\\bar d},s,{\\bar s}$ + gluons &  247 \\\\ $m_{ c} <  T < m_\\tau$ &  $c,{\\bar c}$ & 289 \\\\ $m_\\tau < T < m_{\\textrm{bottom}}$ & $\\tau {}^{\\pm}$ & 303 \\\\ $m_{ b} < T < m_{ W,Z}$ & $b,{\\bar b}$ & 345 \\\\ $m_{ W,Z} <  T < m_{\\textrm{Higgs}}$ & $W^{\\pm}, Z$ & 381 \\\\ $ m_H< T < m_{\\textrm{top}}$ & $H^0$ & 385 \\\\ $m_t< T $ & $t,{\\bar t}$  & 427 \\\\ \\hline \\end{tabular} \\end{center} \\vspace{-.2cm} {\\small *$T_{ c}$ corresponds to the confinement--deconfinement transition between quarks and hadrons. $N(T)$ is shown in Fig.~\\ref{fig:NT} for $T_{c}  =  150$ and $400$~MeV}. \\label{tab:NT} \\end{table}  At higher temperatures, $N(T)$ will be model dependent. For example, in the minimal $SU(5)$ model, one needs to add to $N(T)$, 24 states for the X and Y gauge bosons, another 24 from the adjoint Higgs, and another 6 (in addition to the 4 already counted in $W^\\pm, Z$ and $H$) from the 5 of Higgs. Hence for $T > M_{ X}  $  in minimal $SU(5)$, $N(T)  =  160.75$. In a supersymmetric model this would at least double, with some changes possibly necessary in the table if the lightest supersymmetric particle has a mass below $m_t$.  \\begin{figure} \\centering\\includegraphics[width=.7\\linewidth]{NT} \\caption{The effective numbers of relativistic degrees of freedom as a function of temperature, assuming a quark--hadron transition temperature of 150 MeV and 400 MeV (dashed)} \\label{fig:NT} \\end{figure}   The notion of equilibrium also plays an important role in the standard Big Bang model. If, for example, the Universe were not expanding, then given enough time, each particle state would come into equilibrium with every other.  Because of the expansion of the Universe, certain rates might be too slow indicating, for example, in a scattering process that the two incoming states might never find each other to bring about an interaction. Depending on their rates, certain interactions may pass in and out of thermal equilibrium  during the course of the Universal expansion. Qualitatively, for each particle $i$, we will require that some rate $  \\Gamma {}_{ i} $  involving that type be larger than the expansion rate of the Universe or \\beq \\Gamma {}_{ i}  > H , \\eeq in order to be in thermal equilibrium.  A good example for a  process in equilibrium at some stage and out of equilibrium at others is that of neutrinos. If we consider the standard neutral or charged-current interactions such as $e^{ +} + e^{ -  }     \\leftrightarrow {}   \\nu {}  +  \\bar \\nu $ or $e +   \\nu {}    \\leftrightarrow {}  e  +   \\nu {}$  etc., very roughly the rates for these processes will be \\beq \\Gamma {} =  n \\langle \\sigma v \\rangle , \\eeq where  $  \\langle \\sigma v \\rangle $ is the thermally averaged  weak interaction cross section \\beq \\langle \\sigma v \\rangle { \\sim }~\\mathcal{O}(10^{ -  2}) T^{ 2} /M_{ W}^4 , \\eeq and $n\\sim T^3$ is the number density of leptons. Hence the rate for these interactions is \\beq \\Gamma {}_{\\rm wk}   { \\sim  }~0(10^{ -  2}) T^{ 5}/M_{ W}^4 . \\eeq The expansion rate, on the other hand, is just \\beq H  =  \\left({8 \\pi {}G_{ N}  \\rho {} \\over 3}\\right)^{ 1/2} =  \\left({8  \\pi {}^{ 3} \\over 90}  N(T) \\right)^{ 1/2}  T^{ 2}/M_{ P} \\sim 1.66 N(T)^{ 1/2}  T^{ 2}/M_{ P} . \\eeq The Planck mass $M_{ P} = G_N^{-1/2} =  1.22 \\times 10^{19}$ GeV.  Neutrinos will be in equilibrium when  $  \\Gamma {}_{\\rm wk} >  H $ or \\beq T > (500 M_{ W}^4)/M_{ P})^{ 1/3}  { \\sim }~1 {\\rm ~MeV} . \\eeq The temperature at which these rates are equal is commonly referred to as the decoupling or freeze-out temperature and is defined by \\beq \\Gamma(T_{ d}) = H(T_{ d}) . \\eeq For temperatures $T > T_{ d}$, neutrinos will be in equilibrium, while for $T < T_{ d }$ they will not. Basically, in terms of their interactions, the expansion rate is just too fast and they never {\\em `see'}  the rest of the matter in the Universe (nor themselves). Their momenta will simply redshift and their effective temperature (the shape of their momenta distribution is not changed from that of a blackbody) will simply fall with $T {~\\sim~} 1/R.$  The relation $R T \\sim const$ is a direct consequence of the energy conservation equation~(\\ref{rhod}). Indeed, using $s = dp/dT$, this equation can be rewritten as \\beq \\frac{d}{dt} (R^3 s) = 0 , \\eeq making it more a statement of conservation of (comoving) entropy than energy (which is not conserved in comoving coordinates).  Soon after neutrino decoupling, the $e^{\\pm}$  pairs in the thermal background begin to annihilate (when $T \\la m_e$). Because the neutrinos are decoupled, the energy released heats up the photon background relative to the neutrinos. The change in the photon temperature can easily be computed from entropy conservation. The neutrino entropy must be conserved separately from the entropy of interacting particles. If we denote $T_>$, the temperature of photons, and $e^{\\pm}$ before annihilation, we also have $T_\\nu  = T_>$  as well. The entropy density of the interacting particles at $T = T_>$ is just \\beq s_> =  {4 \\over 3}  {\\rho_> \\over T_>}  = {4 \\over 3} (2 + {7 \\over 2}) ( {\\pi^2 \\over 30} ) T^3_>\t, \\eeq while at $T = T_<$, the temperature of the photons just after $e^{\\pm}$ annihilation, the entropy density is \\beq s_< =  {4 \\over 3}  {\\rho_< \\over T_<}  = {4 \\over 3} (2 ) ( {\\pi^2 \\over 30} ) T^3_<\t, \\eeq and by conservation of entropy $s_< =  s_>$  and \\beq (T_</T_>)^3  = 11/4 . \\eeq Thus, the photon background is at higher temperature than the neutrinos because the $e^{\\pm}$ annihilation energy could not be shared among the neutrinos, and \\beq T_\\nu = (4/11)^{1/3}   T_\\gamma   \\simeq 1.9~{\\rm K} . \\eeq  For further reading on standard Big Bang cosmology see Refs.~\\cite{bbpeeb93,bbborner,bbjap,bbmuk,bbw2,op}.  \\subsection{The Cosmic Microwave Background}  There has been a great deal of progress in the last several years concerning the determination of both $\\Omega_m$ and $\\Omega_\\Lambda$. Cosmic Microwave Background (CMB) anisotropy experiments have been able to determine the curvature (i.e., the sum of $\\Omega_m$ and $\\Omega_\\Lambda$) to better than one per cent, while observations of type Ia supernovae at high redshift  and baryon acoustic oscillations provide information on (nearly) orthogonal combinations of the two density parameters.  The CMB is of course deeply rooted in the development and verification of the Big Bang model and Big Bang Nucleosynthesis (BBN)~\\cite{gamo}. Indeed, it was the formulation of BBN that led to the prediction of the microwave background~\\cite{ah}.  The argument is rather simple. BBN requires temperatures greater than 100 keV, which according to the Standard Model time--temperature relation, Eq.~(\\ref{tT2}), corresponds to timescales less than about 200 s. The typical cross section for the first link in the nucleosynthetic chain is \\begin{equation} \\sigma v (p + n \\rightarrow D + \\gamma) \\simeq 5 \\times 10^{-20} {\\rm ~cm}^3/{\\rm s} . \\end{equation} This implies that it was necessary to achieve a density \\begin{equation} n \\sim {1 \\over \\sigma v t} \\sim 10^{17} {\\rm ~cm}^{-3} , \\end{equation} for nucleosynthesis to begin. The density in baryons today is known approximately from the density of visible matter to be ${n_B}_0 \\sim 10^{-7}$ cm$^{-3}$ and since we know that the density $n$ scales as $R^{-3} \\sim T^3$, the temperature today must be \\begin{equation} T_0 = ({n_B}_0/n)^{1/3} T_{\\rm BBN} \\sim 10 {\\rm ~K} \\end{equation} thus linking two of the most important tests of the Big Bang theory.  Of course it was not until many years later that the microwave background radiation was discovered by Penzias and Wilson~\\cite{pw} while perfecting a radio antenna to track the Echo satellite.  They found a background noise which could not be eliminated corresponding to a temperature of $3.5 \\pm 1$ K. One of the most important papers on modern cosmology was published with the title ``A measurement of excess antenna temperature at 4080-Mc/s''. This was followed by the seminal paper by Dicke, Peebles, Roll, and Wilkinson~\\cite{dprw} putting this observation in a cosmological context. Subsequently, there have been many observations of the CMB culminating in the COBE observation~\\cite{cobe} which determined the temperature to an unprecedented level, set aside any lingering doubts about the true black body nature of the CMB, and discovered the intrinsic anisotropies in the background.  An enormous amount of cosmological information is encoded in the angular expansion of the CMB temperature \\begin{equation} T(\\theta,\\phi) = \\sum_{\\ell m} a_{\\ell m} Y_{\\ell m}(\\theta, \\phi) . \\end{equation} The monopole term characterizes the mean background temperature of $T_\\gamma = 2.725 \\pm 0.001$ K as determined by COBE~\\cite{cobem}, whereas the dipole term can be associated with the Doppler shift produced by our peculiar motion with respect to the CMB. In contrast, the higher order multipoles are directly related to energy density perturbations in the early Universe. When compared with theoretical models, the higher order anisotropies can be used to constrain several key cosmological parameters. In the context of simple adiabatic cold dark matter (CDM) models, there are nine of these: the cold dark matter density, $\\Omega_m h^2$; the baryon density, $\\Omega_B h^2$; the curvature --- characterized by $\\Omega_{\\rm total}$; the hubble parameter, $h$; the optical depth, $\\tau$; the spectral indices of scalar and tensor perturbations, $n_s$ and $n_t$; the ratio of tensor to scalar perturbations, $r$; and the overall amplitude of fluctuations, $Q$.  \\begin{figure} \\centering\\includegraphics[width=.90\\linewidth]{clplotlin_col_09.eps} \\caption{The power in the microwave background anisotropy spectrum as measured by WMAP~\\protect\\cite{wmap}, Boomerang~\\protect{\\cite{boom}}, QUaD~\\protect{\\cite{quad}}, CBI~\\protect\\cite{cbi}, and ACBAR~\\protect\\cite{acbar}. Figure courtesy of D. Scott~\\cite{scsm}. } \\label{freq1} \\end{figure}    Microwave background anisotropy measurements have made tremendous advances in the last few years. The power spectrum~\\cite{max,boom,dasi1,cbi,vsa,arch,wmap0,acbar,quad} has been measured relatively accurately out to multipole moments corresponding to $\\ell \\sim 2000$. A compilation of recent data is shown in Fig. \\ref{freq1}~\\cite{scsm}, where the power at each ${\\ell}$ is given by $(2{\\ell} + 1) C_\\ell/(4 \\pi)$, and $C_\\ell = \\langle |a_{\\ell m}|^2 \\rangle$.       As indicated above, the details of this spectrum enable one to make accurate predictions of a large number of fundamental cosmological parameters. The results of the WMAP data (with other information concerning the power spectrum) is shown in Table~\\ref{tab:wmap}. For details see Refs.~\\cite{wmap,scsm,ll}.    \\begin{table}[h] \\centering \\caption{WMAP determinations of cosmological parameters} \\begin{tabular}{l|cc} & {\\bf WMAP alone} & {\\bf WMAP + BAO + SN} \\\\ \\hline $\\Omega_{{\\rm m}} h^2$ & $0.133 \\pm 0.006$ & $0.136 \\pm 0.004$\\\\ $\\Omega_{{\\rm B}} h^2$ & $0.0227 \\pm 0.0006$ & $0.0227 \\pm 0.0006$\\\\ $h$ & $0.72 \\pm 0.03$ & $0.705 \\pm 0.013$\\\\ $n_s$ & $0.963^{+ 0.014}_{-0.015}$ & $0.960 \\pm 0.013$\\\\ $\\tau$ & $0.087 \\pm 0.017$ & $0.084 \\pm 0.016$\\\\ $\\Omega_\\Lambda$ & $0.74 \\pm 0.03$ & $ 0.726 \\pm 0.015$ \\\\ \\end{tabular} \\label{tab:wmap} \\end{table}      Of particular interest to us here is the CMB determination of the total density $\\Omega_{\\rm tot}$ as well as the matter density $\\Omega_m$. There is strong evidence that the Universe is flat or very close to it. As noted earlier, the best determination of $\\Omega_{\\rm total}$ is $1.006 \\pm 0.006$. Furthermore, the matter density is significantly larger than the baryon density implying the existence of cold dark matter. Also,  the baryon density, as we will see below, is consistent with the BBN production of D/H and its abundance in quasar absorption systems. The apparent discrepancy between the CMB value of $\\Omega_{\\rm tot}$ and $\\Omega_m$, though not conclusive on its own, is a sign that a contribution from the vacuum energy density or cosmological constant, is also required.  The preferred region in the $\\Omega_m - \\Omega_\\Lambda$ plane is shown in Fig.~\\ref{lm}.  \\begin{figure} \\centering\\includegraphics[width=.7\\linewidth]{wmap_snls_bao.eps} \\caption{\\footnotesize Two-dimensional confidence regions in the ($\\Omega_m,\\Omega_{\\Lambda}$) plane. The coloured Monte Carlo points derive from WMAP~\\cite{wmap0} and show that the CMB alone requires a flat Universe $\\Omega_{\\Lambda}+\\Omega_{m}\\simeq 1$ if the Hubble constant is not too high. The SNe Ia results~\\cite{sn2} very nearly constrain the orthogonal combination $\\Omega_{\\Lambda}-\\Omega_{m}$. Also shown is the region obtained from baryon acoustic oscillations~\\cite{kow}. The intersection of these constraints is the most direct (but far from the only) piece of evidence favouring a flat model with $\\Omega_{m}\\simeq 0.25$. } \\label{lm} \\end{figure}    The presence or absence of a cosmological constant is a long standing problem in cosmology.   We know that the cosmological term is at most a factor of a few times larger than the current mass density.  Thus from Eq.~(\\ref{H}), we see that the dimensionless combination $G_N \\Lambda \\la 10^{-121}$. Also shown in Fig.~\\ref{lm}, are the results from SN~Ia~\\cite{sn1,sn2} and baryon acoustic oscillations~\\cite{bao}. Taken together, we are led to a seemingly conclusive picture.  The Universe is nearly flat with $\\Omega_{\\rm tot} \\simeq 1$. However, the density in matter makes up only 23\\% of this total, with the remainder in a cosmological constant or some other form of dark energy. ",
        "conclusions": ""
    },
    "1005/1005.3462_arXiv.txt": {
        "abstract": "We offer a simple parameterization of the rate of star formation in galaxies. In this new approach, we make explicit and decouple the timescales associated (a) with disruptive effects the star formation event itself, from (b) the timescales associated with the cloud assembly and collapse mechanisms leading up to star formation.  The star formation law in near-by galaxies, as measured on sub-kiloparsec scales, has recently been shown by Bigiel et al. to be distinctly non-linear in its dependence on total gas density. Our parameterization of the spatially resolved Schmidt-Sanduleak relation naturally accommodates that dependence.  The parameterized form of the relation is $\\rho_* \\sim \\epsilon\\rho_g/(\\tau_s + \\rho_g^{-n})$, where $\\rho_g$ is the gas density, $\\epsilon$ is the efficiency of converting gas into stars, and $\\rho_g^{-n}$ captures the physics of cloud collapse. Accordingly at high gas densities quiescent star formation is predicted to progress as $\\rho_* \\sim \\rho_g$, while at low gas densities $\\rho_* \\sim \\rho_g^{1+n}$, as is now generally observed.  A variable efficiency in locally converting gas into stars as well as the unknown plane thickness variations from galaxy to galaxy, and radially within a given galaxy, can readily account for the empirical scatter in the observed (surface density rather than volume density) relations, and also plausibly account for the noted upturn in the relation at very high apparent projected column densities. ",
        "introduction": "The Schmidt Law has become widely synonymous with any power-law relation between a (local or even global) star formation rate indicator (be it individual OB stars, HII regions, H$\\alpha$ luminosity, UV luminosity, re-radiated 24$\\mu$m radiation or the strength of a [C IV] cooling line) and a corresponding measure of gas density (originally HI and now more frequently H$_2$, or a summed combination of the two)\\footnote{For a comprehensive and tutorial review see Leroy et al. (2008), Section 2 and especially Table 1.}. Schmidt (1957, 1963) first proposed such a formalism after noting that within the context of the Milky Way the gas scale height (the fuel for star formation) was larger than the O-star scale height (the result of star formation) suggesting a non-linear (plausibly a power-law) causally connected relationship between the two.  Schmidt concluded that the exponent connecting gas density to star formation had a value of n $\\sim$ 2. Subsequently, (with the exception of a solitary paper by Guibery, Lequeux \\& Viallefond 1978, again dealing with the Milky Way star formation and gas scale heights) little interest in this topic was visible for more than a decade.  The field became active again when some years later Sanduleak (1968) offered a novel calibration of the Schmidt Law in an extragalactic context.  Rather than working with scale heights he examined the relationship between neutral hydrogen gas surface densities and the projected surface densities of recently-formed and individually resolved OB stars.  We refer hereafter to this spatially resolved correlation of star formation tracers with projected gas surface density as the Schmidt-Sanduleak Law, so as to clearly distinguish it from the global relation linking the total star formation rate with total gas content, often referred to as the Schmidt-Kennicutt Law. For OB stars and neutral hydrogen gas in the SMC Sanduleak found $\\Sigma_{stars} \\sim \\Sigma^{1.84 \\pm 0.14}_{gas}$.  Sanduleak's particular technique in his pioneering study was extended to other galaxies and in the process also generalized to different star-formation tracers.  For instance, Hartwick (1971) correlated HII region surface densities with neutral hydrogen surface densities for the Local Group galaxy M31, and derived an exponent $n = 3.50\\pm0.12 $. Madore, van den Bergh \\& Rogstad (1974) combined the methods of Sanduleak and Hartwick and correlated both star counts and HII regions with HI surface densities across the face of M33. They found differences\\footnote{And with the clarity of hindsight one can see a flattening of the power-law relation at high gas densities in the M33 data both for the stars and for the HII regions (their Figures 3 and 4); however it was clearly too early in the game to be introducing even more variables than had already been considered in this early paper.}  between between the two tracers as well as differences as a function of radius; the latter, being suggested by them, as being plausibly due to plane thickness variations. They found exponents ranging from $n = 0.57$ to $2.57$, depending on the tracer and the radial region considered, with the outer regions showing a steeper dependence. Tosa \\& Hamajima (1975), Hamajima \\& Tosa, (1975) followed Hartwick's lead and also used HII region number densities and correlated these with the available neutral hydrogen surface densities. They found exponents ranging from $n = 1.5$ to $3.5$ in 7 galaxies (SMC, LMC, M31, M51, M101, NGC~2403 and NGC~6946). Their analysis again suggested that the relation was a function of radius where in the examples of M31 and M101 the exponent was again found to be larger in the outer regions when compared to the inner disk. More recently, Bigiel et al. (2008) examined the spatially resolved star formation law in 18 nearby galaxies using both UV (GALEX) and IR (Spitzer) data to characterize the on-going star formation, and exploring both molecular, neutral and total hydrogen gas content as control variables. They found that for total gas (surface) densities in excess of 1M$_{\\sun}$ pc$^{-2}$ the star formation rate scaled linearly with total gas (i.e., $n = 1.0 \\pm$0.2), but at lower gas surface densities the power law steepens considerably and is widely dispersed (see their composite Figures 10 and 15, and Figure 1 below).  We now proceed to give a physically motivated, analytic expression capturing these facts. It is hoped that with this formulation galaxy evolution modellers will be able to exploit this simple parametric means of predicting star formation rates based on local gas densities. ",
        "conclusions": "The above derivation was made for quiescent star formation, wherein it is implicitly assumed that the rate-limiting timescale behind this particular mode is the coalescence and collapse time of the ambient gas into molecular clouds, which subsequently act as the more immediate sites for star formation. All of the sub-resolution physics and astrophysics of cloud formation and collapse is secreted away in this simple parameterization on the cloud stagnation timescales ($\\tau_s$) and local gas densities ($\\rho^{-n}$). Other modes of star formation of course exist, and because of the very different physics involved, they deserve their own parameterization; but the simple timescale framework should in principle accommodate these other modes as well. The collisionally driven (impulsive) star formation rate $\\tau_c$ would have a timescale set by the local galaxian environment. The internally driven star formation rates due to density waves would have their timescales $\\tau_D$ strictly set locally at a given galactocentric radius by 1/$\\mid \\Omega_{pattern} - \\Omega_{rotation} \\mid$, where the denominator is the rate of passage of material through the density wave of fixed pattern speed. Finally, one could imagine finding an appropriate timescale $\\tau_E$ for the transfer of angular momentum regulating the flow of material into the central regions of galaxies where starburst activity then occurs. In any case the simple dimensionality of the proposed parameterization guarantees some measure of success; the novelty is in decoupling and making explicit the star formation tracer's visibility timescale from the independent assembly/formation/collision timescale.  There are, of course, other proposed parameterizations of star formation rates, many of which are far more bottom-up and physically based rather than the top-down and more phenomenologically motivated approach historically made and adopted here.  Dopita has long argued for a ``compound'' Schmidt Law (e.g. Dopita \\& Ryder 1994 and references therein) that depends both on local gas density and the total surface density of matter.  Silk (1997) argues that the SFR is self-regulated with gravitational instability playing off against supernova heating of the interstellar medium. Tan (2000) suggested that the cloud-cloud collision timescale was the rate-limiting factor for star formation; while Krumhotz \\& McKee (2005) suggested a simple linear relation holds between the surface density of molecular hydrogen and SFR based on the suggestion that the free-fall time for giant molecular clouds is a largely independent of the cloud mass (Solomon et al. 1987). Li, Mac Low \\& Klessen (2005) investigated the role played gravitational instabilities controling star formation rates using a three-dimensional, smooth particle hydro code. Simulations by Dobbs \\& Pringle (2009) simply tie the SFR to the local dynamical timescale of the gas.  Finally, Blitz \\& Rosolowsky (2006 and earlier references therein) suggest that mid-plane pressure is the controling factor in converting HI into its molecular phase and then on to star formation. A detailed comparison of the predictions of some of these models, as well as a number of thresholding scenarios, with the available observations has been given in Leroy et al. (2008) and the interested reader is referred to that paper for a detailed discussion of their relative merits and degrees of success in matching the observations.  The observed rate of star formation is further complicated in the high density nuclear regions of galaxies that are probably not undergoing typical, quiescent star formation that is the main focus of this {\\it Letter}. At very high column densities, often only found in the nuclear regions of galaxies, the raw correlation of star formation rates with very high gas column densities appears to steepen with respect to the relation found the lower-density main disk. Assuming that the molecular fraction has been properly estimated and/or measured, we have two possible explanations for these observations, both of which may be operating in any given situation. First, the plane thinkness in these inner regions may be playing a factor in increasing the volume gas density without necessarily changing the apparent column density. Adopting a standard (outer disk) Schmidt law would result in the observed upturn in the correlation because of the increased volume gas density and the non-linear response in the production of stars each of which are only measured in projection.  In addition to, or independently of, any plane thickness compression at very high (nuclear) gas densities the same steepening of the observed relation can be easily induced by another plausible effect, beam dilution. If, within the naturally imposed or artificially selected resolution element empirically used to make the star-tracer and gas density correlations, the actual size of the region undergoing star formation is smaller than the beam then, as in the case of the decreasing plane thickness situation above, the appropriate volume density of the region actually undergoing star formation will be much higher in reality than the beam-smeared and diluted column density would suggest.  If this is indeed the case, then the observed correlation of star density versus projected gas surface density would steepen, because of the highly non-linear underlying sensitivity of the rate of star formation to volume density.   \\medskip \\medskip"
    },
    "1005/1005.4845_arXiv.txt": {
        "abstract": "{Betelgeuse is an M supergiant with a complex and extended atmosphere, which also harbors spots and giant granules at its surface. A possible magnetic field could contribute to the mass loss and to the heating of the outer atmosphere.} {We observed Betelgeuse, to directly study and infer the nature of its magnetic field.} {We used the new-generation spectropolarimeter NARVAL and the least square deconvolution (LSD) method to detect circular polarization within the photospheric absorption lines of Betelgeuse. } {We have unambiguously detected a weak Stokes $V$ signal in the spectral lines of Betelgeuse, and measured the related surface-averaged longitudinal magnetic field $B_\\ell$ at 6 different epochs over one month.  The detected longitudinal field is about one Gauss and is apparently increasing on the time scale of  our observations.} {This work presents the first direct detection of the magnetic field of Betelgeuse. This magnetic field may be associated to the giant convection cells that could enable a ``local dynamo''.} ",
        "introduction": "Betelgeuse ($\\alpha$ Ori, HD 39801) is a nearby M2Iab supergiant, and a star with the largest apparent diameter. It has naturally deserved direct imaging and interferometric studies of its surface, which was found to deviate from circular symmetry and to present variable behavior (Wilson et al. 1997, Haubois et al. 2009). The presence of a few spots and giant granules appeared as a natural explanation for the imaging and interferometric observations and their variability. A magnetic field at the surface of this star has also been suggested  by Dorch and Freytag (2003). Numerical simulations of convection now match interferometric observations and enable determination of the convection pattern on Betelgeuse (Chiavassa et al. 2010).  In this context, we have undertaken a very sensitive magnetic study of Betelgeuse, using the new generation spectropolarimeter NARVAL in operation at T\\'elescope Bernard Lyot (TBL, Observatoire du Pic du Midi). We present here the result of a short campaign in March-April 2010. We report the definite detection of a weak magnetic field at the surface of Betelgeuse. Section 2 describes our observations and Sect. 3 our results. We discuss the nature of the magnetic field in Sect. 4 and give our conclusions in Sect. 5. ",
        "conclusions": "We Zeeman-detected a magnetic field on the surface of Betelgeuse and measured its surface-averaged longitudinal component to be about 1 G. Because  M supergiants are very slow rotators and since they have very  extended envelopes, neither a solar-type dynamo nor a remnant of a fossil magnetic field from a magnetic main sequence progenitor is expected to be the origin of the magnetic field of Betelgeuse. A possible explanation for its existence can therefore be the giant convective cells predicted to exist on the surface of the star (Schwarzschild 1975), explaining the observations with a high angular resolution (e.g. Haubois et al. 2009, Chiavassa et al. 2010) and appearing in the numerical simulations (Freytag et al. 2002; Dorch 2004). Certainly, the established magnetic field plays a role in the mass loss and in the heating of the outer atmosphere of Betelgeuse (Lim et al. 1998)."
    },
    "1005/1005.4242_arXiv.txt": {
        "abstract": "A time-resolved spectral analysis for a sample of 22 intense, broad GRB pulses from the {\\em CGRO}/BATSE GRB sample is presented. We fit the spectra with the Band function and investigate the correlation between the observed flux ($F$) and the peak energy ($E_{\\rm p}$) of the $\\nu f_\\nu$ spectrum in the rising and decaying phases of these pulses. Two kinds of $E_{\\rm p}$ evolution trends, i.e., hard-to-soft (the two-third pulses in our sample) and $E_{\\rm p}$-tracing-$F$ (the one-third pulses in our sample) are observed in pulses from different GRBs and even from different pulses of the same burst. No dependence of spectral evolution feature on the pulse shape is found. A tight $F-E_{\\rm p}$ positive correlation is observed in the decaying phases, with a power-law index $\\sim 2.2$, which is much shallower than that expectation of the curvature effect. In the rising phase, the observed $F$ is either correlated or anti-correlated with $E_{\\rm p}$, depending on the spectral evolution feature, and the power-law index of the correlation is dramatically different among pulses. More than $80\\%$ of the low energy photon indices in the time-resolved spectra whose $E_{\\rm p}$ is anti-correlated with $F$ during the rising phase violate the death line of the synchrotron radiation, disfavoring the synchrotron radiation model for these gamma-rays. The $F-E_{\\rm p}$ correlation, especially for those GRBs with $E_{\\rm p}$-tracking-$F$ spectral evolution, may be due to the viewing angle and jet structure effects. In this scenario, the observed $F-E_{\\rm p}$ correlation in the rising phase may be due to the line of sight from off-beam to on-beam toward a structured jet (or jitter),  and the decaying phase is contributed by both the on-beam emission and the decayed photons from high latitude of the GRB fireball, resulting in a shallower slope of the observed $F-E_{\\rm p}$ correlation than that predicted by the pure curvature effect. ",
        "introduction": "The physics of prompt emission of gamma-ray bursts (GRBs) remains as a great puzzle. Analysis for a large sample of GRB spectra observed with Burst and Transient Source Experiment (BATSE) on board {\\em CGRO} reveals that the GRB spectra are non-thermal, well-fit with a smoothly-jointed broken power-law, the so-called Band function (Band et al. 1993). The physical radiation mechanism that shapes such a spectrum shape is unclear (e.g., Zhang \\& M\\'{e}sz\\'{a}ros et al. 2004). The peak energy of the $\\nu f_\\nu$ spectrum ($E_{\\rm p}$) is one of the most interesting parameters of GRBs. The relation between the burst energy and $E_{\\rm p}$ may shed light on the physics of GRBs (Zhang \\& M \\'{e}sz\\'{a}ros 2002). With a sample of 12 GRBs detected by BeppoSAX, Amati et al. (2002) found a positive correlation between the isotropic gamma-ray energy $(E_{\\rm iso})$ and the peak energy in the burst frame ($E_{\\rm p,z}$). This correlation was confirmed and even extended to the X-ray flashes by the observations with {\\em HETE}-2 and {\\em Swift} (Sakamoto et al., 2006; Amati 2006). By geometrical-correction for the jet beaming effect or empirically incorporating the breaks in the optical afterglow lightcurves, this correlation is even getting tighter (Ghirlanda et al. 2004; Liang \\& Zhang 2005). Similarly, the isotropic peak luminosity ($L_p$) is correlated with $E_{p,z}$ (Wei et al. 2003; Yonetoku et al. 2004). The average flux ($F$) in a given epoch is also correlated with the $E_{\\rm p}$ in the corresponding epoch within a GRB (Liang et al. 2004). The $F-E_{\\rm p}$ correlation suggests that the $E_{\\rm p}-E_{\\rm iso}$ and the $E_{\\rm p}-L_{p}$ correlations would not be due to the observational selection effects (c.f., Nakar \\& Piran 2005; Band \\& Preece 2005; Shahmoradi \\& Nemiroff 2009). All the prompt GRB emission models predict $E_{\\rm p}$ as a function of both $E_{{\\rm iso}}$ (or $L_{\\rm p}$) and the initial Lorentz factor of the GRB fireball ($\\Gamma_0$) (e.g. Table 1 of Zhang \\& M\\'{e}sz\\'{a}ros 2002). Most recently, Liang et al. (2010) discover a tight correlation between $E_{\\rm iso}$ and $\\Gamma_0$, i.e., $\\Gamma_0=182(E_{\\rm iso}/10^{52}\\rm \\ erg)^{0.25}$. This correlation poses constraints on prompt emission models. For example, the internal shock synchrotron model, the most favorite model for GRBs, predicts $E_{\\rm p} \\propto L^{1/2} \\Gamma_0^{-2}$ (Zhang \\& M\\'{e}sz\\'{a}ros 2002). Combining with the trivial proportionality of $L\\propto E_{\\rm iso}$, one can find that $E_{\\rm p}$ should not depend on $L$, indicating that the $E_{\\rm p}-E_{\\rm iso}$ correlation may not be explained in the framework of the internal shock synchrotron model.  This paper dedicates to revisit the $F-E_{\\rm p}$ correlation and its possible physical origin. The $E_{\\rm p}$ of a given burst evolves with time, tracing with the lightcurve ($E_{\\rm p}$-tracing-$F$) or evolving as hard-to-soft (Liang \\& Kargatis 1996; Ford et al. 1995; Kaneko et al. 2006). Intuitively, the $F-E_{\\rm p}$ correlation should be the foundation of the Amati relation, if both are based on the same physical origin. Since those GRBs having hard-to-soft spectral evolution do not have a coherent $F-E_{\\rm p}$ correlation, they should violate the Amati relation. However, the current sample of GRBs with both $E_{\\rm iso}$ and $E_{\\rm p}$ measurements well follow the Amati-relation (Amati et al. 2009), regardless of the spectral evolution feature of these GRBs. For example, although the spectrum of GRB 060218 evolves as hard-to-soft (Toma et al. 2007; Dong et al. 2010), it well satisfies the Amati-relation (Amati et al. 2006). Therefore, the Amati relation seems to be unrelated to the dependence of $F$ on $E_p$. This gives rise to a big puzzle of the physical origin of the Amati-relation.  It is well known that the GRB lightcurvs are generally composed of some overlapped pulses. An individual shock episode may give rise to a pulse, and the random superposition of many such pulses results in the observed complexity of GRB light curves. As the building blocks of GRB lightcurves, the broad, well-separated pulses are good candidates to reveal the physics of GRBs. The $E_{\\rm p}$ evolution within GRB pulses has been extensively studied, especially for the decaying phase of pulses (Golenetskii et al. 1983; Norris et al. 1986; Kargatis et al. 1994, 1995; Bhat et al. 1994; Crider et al. 1999; Peng et al. 2009; Lu \\& Liang 2010). While the decaying phase of a GRB pulse might be contributed by the curvature effect of the GRB fireball (e.g., Fenimore et al. 1995; Kobayashi et al. 1997; Qin et al. 2002; Dermer 2004; Shen et al. 2005), both the temporal and spectral behaviors of the rise phase of a pulse may depend on the dynamics of the GRB fireball, the electron acceleration, and the radiation mechanism (e.g., Kobayashi et al. 1997). The $F-E_{\\rm p}$ correlations in the two phases, if any, should dramatically different, which was shown by Ohno et al. 2009 and Ghirlanda et al. (2010) for some bright GRBs. We here present a detailed analysis on the $F-E_{\\rm p}$ relation in both the rising and decaying phases of smooth, broad GRB pulses observed with {\\em CGRO}/BATSE.  We present our sample and spectral analysis in Section 2. The lightcurves and the $E_p$ evolution as well as the $F-E_{\\rm p}$ correlation are shown in Section 3. Since the low energy photon index of the Band function is critical to justify if the radiation is from synchrotron radiation, we present the distribution of the low energy photon indices for those spectra having an anti-correlation between $E_{\\rm p}$ and $F$ in Section 4. We discuss the possible physical origin of the $F-E_{\\rm p}$ correlation (or anti-correlation ) in Section 5. Conclusions are presented in Section 6. ",
        "conclusions": "With our time-resolved spectral analysis for a sample of 22 intense, broad GRB pulses we find that the $E_{\\rm p}$ evolution feature is well classified into two groups, i.e., hard-to-soft (two-third pulses in our sample) and $E_{\\rm p}$ tracing-$F$ (one-third pulses in our sample). Two kinds of spectral evolutionary trends are also observed in different pulses of a burst. No dependence of spectral evolution feature on the pulse shape is observed.  A tight $F-E_{\\rm p}$ correlation, $F\\propto E^{\\kappa_d}$, is observed in the decaying phases of these pulse. Although the $\\kappa_{d}$ ranges in a broad range, from $0.6$ to $ \\sim 4.0$, their distribution normally peaks at $\\sim 2$, much shallower than that expectation of the curvature effect. In the rising phase, the observed $F$ is either correlated or anti-correlated with $E_{\\rm p}$, $F\\propto E_{\\rm p}^{\\kappa_r}$, depending on the spectral evolution feature. The distribution of $\\kappa_r$ spans from $-4$ to 10 without any correlation with the pulse shape and $\\kappa_d$. More than $80\\%$ of the low energy photon indices in the time-resolved spectra whose $E_{\\rm p}$ is anti-correlated with $F$ violate the death line of the synchrotron radiation.  The spectral evolution features and the observed $F-E_{\\rm p}$ correlation are difficult to explain with the radiation models. We propose that the observed the $F-E_{\\rm p}$ correlation observed in a given pulse, especially  for those GRBs with $E_{\\rm p}$-tracking-$F$ spectral evolution, may be due to the viewing angle and jet structure effects. In this scenario, the observed $F-E_{\\rm p}$ correlation in the rising phase is due to the line of sight from off-beam to on-beam toward a structured jet. The contribution of the high latitude photons from the fireball(the so-called curvature effect) would be increased in the decaying phase of a pulse and it should dominate the observed flux when the line of sight moves out the jet edge. Therefore, the mixed contributions from the on-beam and the curvature effect may result in a shallower decay slope of the observed $F-E_{\\rm p}$ correlation as theoretical prediction."
    },
    "1005/1005.1651_arXiv.txt": {
        "abstract": "Intracluster light (ICL) is believed to originate from the stars stripped from cluster galaxies. They are no longer gravitationally bound to individual galaxies, but to the cluster, and their smooth distribution potentially makes them serve as much denser tracers of the cluster dark matter than the sparsely distributed cluster galaxies. We present our study of the ICL in the galaxy cluster Cl~0024+17 using both Advanced Camera for Surveys (ACS) and Subaru data, where we previously reported discovery of a ring-like dark matter structure with gravitational lensing. The ACS images provide much lower sky levels than ground-based data, and enable us to measure relative variation of surface brightness reliably. This analysis is repeated with the Subaru images to examine if consistent features are recovered despite different reduction scheme and instrumental characteristics. We find that the intracluster light profile clearly resembles the peculiar mass profile, which stops decreasing at $r\\sim50\\arcsec$ ($\\sim265$ kpc) and slowly increases until it turns over at $r\\sim75\\arcsec$ ($\\sim397$ kpc). This feature is seen in both ACS and Subaru images for nearly all available passband images while the features are in general stronger in red filters. The consistency across different filters and instruments strongly rules out the possibility that the feature might come from any residual, uncorrected calibration errors. In addition, our re-analysis of the cluster X-ray data shows that the peculiar mass structure is also indicated by a non-negligible (3.7~$\\sigma$ in $Chandra$ and 2.4~$\\sigma$ in XMM-$Newton$) bump in the intracluster gas profile when the geometric center of the dark matter ring, {\\it not the peak of the X-ray emission}, is chosen as the center of the radial bin. The location of the gas ring is closer to the center by $\\sim15\\arcsec$ ($\\sim80$ kpc), raising an interesting possibility that the ring-like structure is expanding and the gas ring is lagging behind perhaps because of the ram pressure {\\it if} both features in mass and gas share the same dynamical origin. ",
        "introduction": "} Intracluster light (ICL), which was discovered by Zwicky (1951), is believed to originate from the stars stripped from cluster galaxies. Although details are still in dispute, a number of studies suggest that both the scatter of stars during the brightest cluster galaxy formation (e.g, Gerhard et al. 2007) and the tidal disruption of dwarf galaxies (e.g., Mihos et al. 2005) are among the dominant mechanisms of the ICL production. These stars are in general considered bound not to any individual galaxy, but to the cluster mass. If a large fraction of the ICL is produced during the assembly of the brightest cluster galaxies (e.g., Murante et al. 2004), the dynamically collisionless property of the intracluster stars potentially allows us to use ICL as visible tracers of underlying dark matter at least in the central region of the cluster.  Observationally, however, the quantitative study of ICL is difficult. The typical surface brightness of ICL is often quoted as $\\sim1$\\% or less of the night sky from the ground. At this faint level, instrumental systematic effects such as residual flat-fielding, scattered lights of bright stars, etc. become critical issues. In addition, most galaxy edges continuously blend into ICL, which obviously causes an ambiguity in determining where the galaxy light stops and ICL begins.  In this paper, we present a study of ICL in the galaxy cluster Cl 0024+17 at $z=0.4$, where we recently discovered a peculiar ringlike dark matter structure (Jee et al. 2007; J07 hereafter). The two-dimensional mass reconstruction of J07 shows that the core of the cluster is surrounded by a $r\\sim0.4$ Mpc ring-like overdense region. The feature is strongly constrained by coherent fluctuation of background galaxy shapes across the $r\\sim0.4$ Mpc circle (weak-lensing) signaling the sudden change of the density slope. Of course, it is tempting to try to detect the peculiar dark matter structure with the cluster galaxies (e.g., Qin et al. 2008). However, because the density contrast of the feature with respect to the neighboring region is low ($<5$\\%), the sparsely distributed galaxies cannot provide sufficient statistics to overcome the shot noise even if the cluster galaxy membership identification is next to perfect. This is the reason that here we investigate, instead, the intracluster stars, which diffusely distribute in the cluster potential and thus sample the underlying dark matter much more densely. In this investigation, the measurement accuracy is not limited by the poissonian scatter as in the study using galaxies, but by the ability to control various systematics critical to precision surface photometry.  We employ two sets of data: {\\it Hubble Space Telescope}  Advanced Camera for Surveys (HST/ACS) and Subaru/Suprime Cam images. Space-based imaging provides significant advantage over ground-based effort mainly because the sky is substantially darker (e.g., $\\sim1$ and $\\sim3$ magnitude fainter in $r$ and $z$ band, respectively). Also, as the instrument is above the atmosphere, there is no cumbersome, time-dependent airglow effect, which often prevents stable calibration. However, unfortunately, there have been no ICL studies so far with the HST/ACS data. This is because the pipeline flatfielding accuracy $\\sim1$\\% (Mack et al. 2002) has been considered insufficient to enable such studies. Therefore, we undertake this time-consuming task of producing/verifying HST/WFC flats using ``blank images'' before we proceed with the ICL measurements. One drawback in using ACS data is the small field of view of the instrument, which only covers a $3\\arcmin\\times3\\arcmin$ region. Because the existing ACS data of the cluster were taken without any plan for this kind of study, the instrument pointing was almost fixed to the center of the cluster (with only a few tens of pixels dither to fill the two CCD gap), and thus a care must be taken to minimize the ICL contamination in determining the background level.  Although the sky is much brighter, the Subaru/Suprime Cam images provide a few critical crosschecks. First, the large field of view of the Suprime Cam allows us to estimate the background level directly, which is important in determining the net amount of ICL and assessing the degree of the aforementioned ICL contamination in the HST/ACS measurements. Second, the Subaru data include the NB$_{912}$ narrow band image, which Kodama et al. (2004) used to probe H$\\alpha$ emission as a measure of on-going star formation of the cluster at $z=0.4$. The sky level in this narrow band is about 1.5 mag darker than in the broadband $z^{\\prime}$ filter while the narrow passband gives high contrast to the H$\\alpha$ emission at $z=0.4$ relative to continuum sources at different redshifts. Third, various instrumental signatures including the point spread function (PSF), internal reflection, obscuration, geometric distortion, etc. are different from ACS. Hence, if a significant feature clearly revealed in one instrument, is absent in the other, this indicates that the feature might come from some residual calibration errors or time-dependent sources. On the contrary, detection of a consistent feature in both images provides evidence against these systematics.  The hot intracluster gas also samples the underlying dark matter very densely although its collisional nature biases the distribution from that of dark matter. A common isothermal $\\beta$ model assumes $\\rho_{gas} \\propto \\rho_{DM}^{\\beta}$ and $\\epsilon_{ff} \\propto n_e^2$, where $\\epsilon_{ff}$, and $n_e$ are the free-free radiation emissivity and the electron density, respectively. Therefore, X-ray profiles in general are more centrally peaked than mass profiles. Also, as observed in many merging systems (Jee et al. 2005; Clowe et al. 2006), some gas peaks are often significantly displaced with respect to dark matter peaks, which is interpreted as signaling the different physical property (e.g., collisional versus collisionless nature). Therefore, a rigorous simulation is desired to fully address whether or not it is feasible to detect the analogous ring-like structure in the gas profile of Cl~0024+17, given the low level of the contrast in the dark matter structure. In this paper, we approach the issue purely observationally by re-analyzing the existing $Chandra$ and XMM-$Newton$ X-ray data. The current analysis is different from J07 only in that now we choose the geometric center of the dark matter ring as our origin in setting up radial bins. If there is indeed any feature in the gas profile hinting at the peculiar dark matter substructure, the signal is expected to be of low contrast and thus should be sensitive to the choice of the origin. Therefore, it is quite possible that J07 might have failed to detect the feature in X-ray because the center was not placed at the optimal location.  In this paper, we adopt a $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.3$, and $h=0.7$ cosmology, where the plate scale is 5.3 kpc$/\\arcsec$ at the redshift of the cluster $z=0.4$. The quoted uncertainties are at the $1-\\sigma$ ($\\sim68$\\%) level. ",
        "conclusions": "Our main results of the investigation of the ICL and ICM in the galaxy cluster Cl~0024+17 are summarized as follows.  \\begin{itemize} \\item The ICL profile closely resembles the peculiar dark matter structure reported in J07, which stops decreasing at $r\\sim50\\arcsec$ ($\\sim265$ kpc) and slowly increases until it turns over at $r\\sim75\\arcsec$ ($\\sim397$ kpc). \\item The feature is present in both ACS and Subaru images in nearly all available passband images with a significance in general higher in red filters. \\item The ICM profile shows a bump at $r\\sim60\\arcsec$ ($\\sim318$ kpc). The radius of the ICM ring is $\\sim20$\\% smaller than that of the dark matter ring. \\item The two dimensional map of the diffuse light in the cluster shows that the ICL distribution is elongated toward Northwest following cluster members. \\item The ICL color is consistent with the cluster red-sequence color at small radii ($r<40\\arcsec$), but becomes bluer for increasing radii. \\item The ICL fraction of Cl~0024+17 is $\\sim28$\\% ($\\sim37$\\%) in F775W and F850LP (F475W, F555W, F625W) without any strong indication of radial dependence. Considering the redshift of the cluster, these values are likely to be at the high end of the distribution. \\end{itemize}   These results not only support the presence of the peculiar dark matter structure of J07 revealed by coherent fluctuation of the background galaxy shapes across the $r\\sim75\\arcsec$ circle, but also demonstrates the usefulness of the ICL as a powerful probe of dark matter substructure in galaxy clusters.  Theoretically, however, it is still disputable whether or not the ICL distribution for a $z=0.4$ cluster should reveal the underlying dark matter. Numerical simulations by Rudick et al. (2006) show that the ICL morphology is a strong function of time. At early epoch most intracluster stars reside in outer halos of individual galaxies while with time more tidal and filamentary structures develop. As the evolution progresses further, it is shown that these transitory features disappear into the more diffuse, common cluster halo. If we associate the redshift $z=0.4$ with the young dynamical age in ICL evolution, it is difficult to understand why the cluster's ICL follows the dark matter structure so remarkably well as shown in this work. On the other hand, if Cl~0024+17 is regarded dynamically already mature based on its optical richness or mass, the similarity between the ICL and dark matter structures may not be surprising. The high ICL fraction of the cluster supports the possibility that Cl~0024+17 is already mature in the context of the ICL evolution. Alternatively, it is also plausible that the high-speed collision 1-2 Gyrs ago might have liberated unusual amount of stars from the cluster galaxies.  The detection of the ICM bump at $r\\sim60\\arcsec$ opens an important opportunity to study the dynamical interplay between the intracluster medium and the underlying dark matter. One simplistic scenario that one can make out of this feature is that both the ICM and dark matter ringlike structures have the same dynamical origin, and both rings (or shells) are expanding. In this case, the expansion of the ICM ring must lag behind the dark matter ring because of the ram pressure, which may explain the observed $\\sim15\\arcsec$ ($\\sim80$ kpc) offset.  In a recent dark matter only simulation, ZuHone et al. (2008) claim that a bump can arise in cluster collisions only when the initial particle velocity distribution is circular, and the feature becomes a ``shoulder\" as the tangential anisotropy decreases. Our re-examination of the J07 collision simulation reveals that the initial orbits of the particles in the cluster core were indeed biased toward tangential anisotropy although the velocity anisotropy is not as extreme (i.e., circular) as shown by ZuHone et al. (2008). As the distinction between ``bump\" and ``shoulder\" is whether or not the slope of the post-collision mass profile in the intermediate range (i.e.,$50\\arcsec<r<75\\arcsec$) is positive, we believe that many other parameters such as the slope of the initial density profile, the fraction of the gas particles, the relative mass ratio/speed of the two subclusters, etc. should be also considered together with the initial distribution of the particle orbits in order to fully address the issue. Nevertheless, it is clear from the ZuHone et al. (2008) simulation that the creation of the ringlike structure, if the high-speed collision is to be the cause as suggested by J07, requires some degree of tangential anisotropy at least in the core of the cluster.  We acknowledge Maxim Markevitch for first informing us of the presence of the ringlike structure in the Chandra image of Cl 0024+17. We would like to thank Bernard Fort for suggesting this study at the 23rd IAP colloquium. Furthermore, we would like to thank Tony Tyson, Holland Ford, David Wittman, and John Blakeslee for useful discussions. This research is in part supported by the grant from the TABASGO foundation awarded in the form of the Large Synoptic Survey Telescope Cosmology Fellowship.  \\clearpage   \\begin{figure} \\plotone{ICL_illustrate.eps} \\caption{ACS and Subaru color images of Cl~0024+17 and two-dimensional ICL map. The top panel shows the central $10\\arcmin\\times10\\arcmin$ region of the color-composite Subaru image. White circles denote spectroscopic cluster members ($0.37<z<0.42$). The yellow outline illustrates the $3\\arcmin\\times3\\arcmin$ region observed with ACS. The ACS color-composite image is shown separately in the lower left panel. The two red dashed-circles delineate the $r\\sim40\\arcsec$ and $r\\sim70\\arcsec$ radii, across which a sudden change of slope in the ICL profile occurs. The lower right panel displays the diffuse light in ACS F625W image. This two-dimensional ICL image is created by masking out astronomical objects, subtracting bright stars, and then applying $3\\arcsec\\times3\\arcsec$ box median-smoothing. The color scale is linear. The diffuse light extends towards the northwestern substructure indicated by the cluster members and the mass reconstruction of Kneib et al. (2003). \\label{fig_ICL_2d} } \\end{figure}     \\begin{figure} \\includegraphics[width=8cm]{flat_all.eps} \\caption{Residual sky flats in the six ACS/WFC broadband filters. These flats are created from the pipeline-reduced FLT files, and should reveal any residual features that are not corrected by the pipeline LP-flats (see text for details in the procedure). The displayed gray scale is linear. The rms residual is $\\sim0.5$\\%. The most outstanding large scale feature is the donut-like pattern particularly clear in F775W and F850LP. This residual feature may arise from the limited accuracy of the polynomial interpolation in the L-flat creation or from the difference of the spectra between the sky and the 47 Tuc stars as suggested by Pavlovsky et al. (2005). One of the most conspicuous small scale features is the gridlike pattern. We verified that this is not an artifact of the 32$\\times$32 box median smoothing that we apply before stacking is carried out. Although we suspect that the feature might be associated with the CCD fabrication procedure, the exact cause of this $\\sim60$ pixel-scale feature is still under investigation. In addition, the residual flats reveal some new dust-moats and under-sensitive CCD columns. \\label{fig_flat_all} } \\end{figure}    \\begin{figure} \\plotone{masking_effect.eps} \\caption{Experiment with different masking sizes. We investigated the effect of the masking size on the centroids of the Gaussian curves by varying masking sizes and repeating the measurements. As an example, shown here is the result from the $r\\sim80\\arcsec$ bin in F775W. The different Gaussian curves represent the results from different masking schemes. The vertical dashed lines show the corresponding centroids. We observe that the centroid converges when the major and minor axes of the masking ellipse are three times or greater than SExtractor measurements. We conservatively selected five times the values given by SExtractor in defining our masking ellipses throughout the analysis. \\label{fig_masking} } \\end{figure}   \\begin{figure} \\plottwo{ACS_psf.eps}{Subaru_psf.eps} \\caption{Extended PSF wings of ACS/WFC (left) and Subaru/Suprime-cam (right). The normalization is chosen in an arbitrary way so that the profiles from different filters approximately overlap at $r<20\\arcsec$. The ACS results are obtained from the archival observation of the star HD39060 (Beta Pictoris). The Subaru PSFs were measured from the brightest star ($r\\sim8.6$ mag) in the Cl~0024+17 image ($\\sim7\\arcmin$) away from the cluster center. \\label{fig_psf_wing} } \\end{figure}      \\begin{figure} \\plotone{pixel_hist.eps} \\caption{Pixel intensity distribution in the ACS/WFC F775W image. Objects were masked and bright stars were subtracted before the statistics was evaluated. Each curve represents the pixel intensity distribution in each radial bin. The bottom curve represents the data in the $r=8\\arcsec$ bin. To improve readability, we displace the histograms vertically. The top curve corresponds to the data in the $r=104\\arcsec$ bin. We adopt the center of the Gaussian peak as the ICL level of the bin (filled square and thick solid line). The uncertainties in determining the centroid of the Gaussian curves after including photon noise and flatfielding errors are $\\sim0.5$ counts, less than the size of the filled square symbols. \\label{fig_many_histos} } \\end{figure}    \\begin{figure} \\plottwo{icl.eps}{icl_subaru.eps} \\caption{ICL profile in Cl~0024+17. The thick solid lines represent the surface brightness while the dashed (dotted) lines show the 1-$\\sigma$ limits without (with) errors in background level determination included. In case of ACS, we arbitrarily assume that the error in background level measurement is at the 0.1\\% of the sky level. Because of the limited field of view, ICL measurement from the ACS data is obtained from an incomplete circle at $r\\gtrsim85\\arcsec$ (dot-dashed). \\label{fig_icl} } \\end{figure}   \\begin{figure} \\plotone{masstoicl.eps} \\caption{Comparison of the ICL profiles with the mass density profile. The ICL profiles from nearly all the filters reveal the critical features of the mass profile: the steep decline at $r\\lesssim50\\arcsec$, the slow rise at $50\\arcsec \\lesssim r \\lesssim 75\\arcsec$, and the turnaround at $r\\sim75\\arcsec$. However, we note that the ICL profiles decrease faster after their turnover at $r\\simeq72\\arcsec$ than the mass profile of J07. We display the NFW fit result (obtained by forcing the profile to match the mass at $20\\arcsec<r<40\\arcsec$) in order to illustrate how much both the cluster mass and ICL profiles deviate from this analytic model frequently used to describe relaxed clusters. \\label{fig_masstoicl} } \\end{figure}   \\begin{figure} \\includegraphics[width=8cm]{quadrant.eps} \\caption{Same as in Figure~\\ref{fig_icl} except here we measure the ICL profile in each quadrant of the ACS F775W image. Despite the quadrant-to-quadrant variations, the signal of the peculiar ICL structure is present in all four quadrants. The dashed-line does not include the uncertainty from background measurement. \\label{fig_icl_f775w} } \\end{figure}   \\begin{figure} \\plottwo{cm.eps}{icl_color.eps} \\caption{Galaxy color versus ICL color. (a) Color magnitude diagram in the ACS field. We used SExtractor's isophotal magnitude to compute the galaxy colors while the F625W magnitude in the X-axis is SExtractor's MAG\\_AUTO. The red solid line approximately indicates the location of the cluster red-sequence. (b) ICL color profile as a function of radius. The dotted lines represent the 1-$\\sigma$ uncertainty including errors in photon statistics, flatfielding inaccuracy, and background level determination. We assume that the error in the background level determination is at the 0.1\\% level of the sky. \\label{fig_icl_color} } \\end{figure}   \\begin{figure} \\includegraphics[width=8cm]{icl_fraction.eps} \\caption{Ratio of ICL flux to total cluster flux. We combined the photometric redshift catalog of J07, and the publicly available spectroscopic catalog of Moran et al. (2005) to identify the cluster members. \\label{fig_icl_fraction} } \\end{figure}     \\begin{figure} \\plotone{xray_profile.eps} \\caption{X-ray surface brightness profile created from the $Chandra$ data. (a) When the origin agrees with the center of the dark matter ring, it reveals an overdensity at $r\\sim60\"$ (the inset picture shows the point-source removed X-ray image and the `+' symbol denotes the origin of the radial bin). (b) When we align the origin to the peak of the X-ray emission (also the location of the mass and luminosity peak), the bump in (a) becomes much less prominent. We let the neighboring bins overlap by $\\sim50$\\% to suppress the shot noise. \\label{fig_xray_profile} } \\end{figure}  \\begin{figure} \\plotone{xmm-newton.eps} \\caption{X-ray surface brightness profile from the XMM-$Newton$ data. The bump at $r\\sim60\\arcsec$ seen in Figure~\\ref{fig_xray_profile} is somewhat smoothed by the larger PSF but still visible as a shoulder. The feature is difficult to observe if we chose the X-ray peak as the origin of the radial bin as in the case of Figure~\\ref{fig_xray_profile}. The scale is arbitrary. The neighboring bins do not overlap. \\label{fig_xray_profile_xmmnewton} } \\end{figure}   \\begin{figure} \\plotone{mass_compare.eps} \\caption{Discrepant mass profiles of Cl~0024+17 in the literature.  The unit is the critical lensing mass density at the $z=3$ reference redshift.  For the Broadhurst et al. (2000) results, we reproduce their two-dimensional mass map, and measure the profile by placing the center at the geometric center of the ringlike structure of J07.  For the Tyson et al. (2000) and Kneib et al. (2003) we use their parametric fit to the mass profiles. The Zitrin et al. (2009) profile shown here is their best-fitting model (read off from Figure 5 of the paper). \\label{fig_various_mass_profiles} } \\end{figure}   \\begin{figure} \\plotone{galaxy_vs_mass.eps} \\caption{Projected galaxy number density versus radius. We use the publicly available spectroscopic catalog of Moran et al. (2005) to plot the number density of the spectroscopic members in the cluster as a function of radius (filled circle). Because the mass profile of J07 is progressively biased beyond $r\\gtrsim80\\arcsec$ (dashed line) because of the limited field of view, a fair comparison may require limiting the range to fit at $r<80\\arcsec$. This gives a goodness-of-the-fit of $\\chi^2/d.o.f=0.45$ (solid), which illustrates that the galaxy distribution is statistically consistent with the J07 mass profile. If the data points at $r>80\\arcsec$ are included, the goodness-of-the-fit degrades to $\\chi^2/d.o.f=1.92$ (dotted). However, interpretation needs a caution because the result is severely affected by the data points at $r>80\\arcsec$, which have smaller error bars yet with lower completeness. \\label{fig_galaxy_vs_mass} } \\end{figure}   \\clearpage  \\begin{deluxetable}{ccccc} \\tabletypesize{\\scriptsize}  \\tablecaption{ACS Image Sky Levels}  \\tablenum{1}  \\tablehead{\\colhead{Filter} & \\colhead {Visits} & \\colhead{Total Exp. Time} & \\colhead{Mean Sky Level} & \\colhead{Sky Level Variation Between Visits} \\\\ \\colhead{} &  \\colhead{}      & \\colhead{ (s) }     & \\colhead{ (mag~$\\mbox{arcsec}^{-2}$)} & \\colhead{ (\\%)} } \\tablewidth{0pt}  \\startdata F435W  & 5 & 6435  &  23.1 & 11\\\\ F475W  & 4 & 5072  &  22.9 & 4 \\\\ F555W  & 4 & 5072  &  22.6 & 3 \\\\ F625W  & 7 & 8971  &  22.2 & 8 \\\\ F775W  & 8 & 10144 &  22.2 & 2 \\\\ F850LP & 6 & 16328 &  22.1 & 2 \\\\ \\enddata \\end{deluxetable}  \\begin{deluxetable}{ccccc} \\tabletypesize{\\scriptsize}  \\tablecaption{Subaru Image Sky Levels}  \\tablenum{2}  \\tablehead{\\colhead{Filter} & \\colhead{Visits} & \\colhead{Total Exp. Time} & \\colhead{Mean Sky Level at zenith} & \\colhead{Sky Level Variation Between Visits} \\\\ \\colhead{}  & \\colhead{}     & \\colhead{ (s) }     & \\colhead{ (mag~$\\mbox{arcsec}^{-2}$)} & \\colhead{ (\\%)}  } \\tablewidth{0pt}  \\startdata $B$           & 3  & 3,600  &  22.6 & 2  \\\\ $R_c$         & 11 & 5,280  &  21.2 & 3  \\\\ $z^{\\prime}$  & 8  & 1,980  &  19.2 & 5  \\\\ $NB_{912}$      & 6  & 10,800 &  20.7 & 11 \\\\  \\enddata \\end{deluxetable}   \\begin{deluxetable}{ccccc} \\tabletypesize{\\scriptsize} \\tablecaption{ACS Sky Flat Accuracy} \\tablenum{3}  \\tablehead{\\colhead{Filter} & \\colhead{Number of Used Images} & \\colhead{Deviation from Pipeline Flats} & \\colhead{Accuracy} \\\\ \\colhead{}  & \\colhead{} & \\colhead{ (\\%)} & \\colhead{ (\\%)} } \\tablewidth{0pt} \\startdata F435W  &90  &  0.5  &  0.08  \\\\ F475W  &138 &  0.4  &  0.07  \\\\ F555W  &92  &  0.4  &  0.06  \\\\ F625W  &155 &  0.5  &  0.07  \\\\ F775W  &154 &  0.4  &  0.07  \\\\ F850LP &312 &  0.7  &  0.08  \\\\ \\enddata \\end{deluxetable}    \\begin{deluxetable}{ccc} \\tabletypesize{\\scriptsize} \\tablecaption{Subaru Sky Flat Accuracy} \\tablenum{4} \\tablehead{\\colhead{Filter} & \\colhead{Number of Used Images} & \\colhead{Accuracy}   \\\\ \\colhead{}     & \\colhead{  }     & \\colhead{ (\\%)} } \\tablewidth{0pt} \\startdata $B$            &49  &   0.08 \\\\ $R_c$          &138 &   0.05 \\\\ $z^{\\prime}$   &155 &   0.05 \\\\ $NB_{912}$     &41 &    0.09 \\\\ \\enddata \\end{deluxetable}    \\begin{deluxetable}{cccc} \\tabletypesize{\\scriptsize} \\tablecaption{Errors in Subaru Background Level Estimation} \\tablenum{5} \\tablehead{\\colhead{Filter} & \\colhead{Fraction of Sky} & \\colhead{Surface Brightness}  \\\\ \\colhead{}     & \\colhead{ (\\%)  }     & \\colhead{ ($\\mbox{mag}~\\mbox{arcsec}^{-2}$) }} \\tablewidth{0pt} \\startdata $B$          &    0.05 & 30.9\\\\ $R_c$        &    0.04 & 29.7\\\\ $z^{\\prime}$ &    0.04 & 27.6\\\\ $NB_{912}$   &    0.08 & 28.4\\\\ \\enddata \\end{deluxetable}   \\begin{deluxetable}{cccc} \\tabletypesize{\\scriptsize} \\tablecaption{ICL fraction in Cl 0024+17} \\tablenum{6} \\tablehead{\\colhead{Filter} & \\colhead{$r<200$ kpc} & \\colhead{$r<500$ kpc}  \\\\ \\colhead{} & \\colhead{(\\%)} & \\colhead{(\\%)} } \\tablewidth{0pt} \\startdata F475W  &    $35\\pm6$  &    $38\\pm12$\\\\ F555W  &    $35\\pm4$   &    $37\\pm9$\\\\ F625W  &    $35\\pm2$   &    $35\\pm5$\\\\ F775W  &    $27\\pm1$  &    $26\\pm3$\\\\ F850LP  &   $29\\pm1$   &    $30\\pm3$\\\\ \\enddata \\end{deluxetable}  \\clearpage"
    },
    "1005/1005.5244_arXiv.txt": {
        "abstract": "We present a multiwavelength analysis of a long duration white-light solar flare (M8.9/3B) event that occurred on 04 June 2007 from AR NOAA 10960. The flare was observed by several spaceborne instruments, namely SOHO/MDI, HINODE/SOT, TRACE and STEREO/SECCHI. The flare was initiated near a small, positive-polarity, satellite sunspot at the centre of the active region, surrounded by opposite-polarity field regions. MDI images of the active region show considerable amount of changes in the small positive-polarity sunspot of $\\delta$ configuration during the flare event. SOT/G-band (4305 \\AA) images of the sunspot also suggest the rapid evolution of this positive-polarity sunspot with highly twisted penumbral filaments before the flare event, which were oriented in a counterclockwise direction. It shows the change in orientation, and also remarkable disapperance of twisted penumbral filaments ($\\approx$35\\,--\\,40$\\%$) and enhancement in umbral area ($\\approx$45\\,--50\\,$\\%$) during the decay phase of the flare event. TRACE and SECCHI observations reveal the successive activation of two helical-twisted structures associated with this sunspot, and the corresponding brightening in the chromosphere as observed by the time-sequence images of SOT/Ca\\,{\\sc ii} H line (3968 \\AA). The secondary, helical-twisted structure is found to be associated with the M8.9 flare event. The brightening starts \\,six--seven\\, minutes prior to the flare maximum with the appearance of secondary, helical-twisted structure. The flare intensity maximizes as the secondary, helical-twisted structure moves away from the active region. This twisted flux tube, associated with the flare triggering, is found to be failed in eruption. The location of the flare activity is found to coincide with the activation site of the helical twisted structures. We conclude that the activation of successive helical twists (especially the second one) in the magnetic flux tubes/ropes plays a crucial role in the energy build-up process and triggering of the M-class solar flare without a coronal mass ejection (CME). ",
        "introduction": "Solar flares are the sudden explosions in the solar atmosphere during which large amounts of the magnetic energy, stored in the twisted and sheared magnetic fields, is released by the process of magnetic reconnection in the form of thermal energy and particle acceleration. The flares associated with the CME eruptions are known as ``eruptive flares'', while the flares without association of CMEs are known as ``confined flares''. The emerging magnetic flux,  rapid  motion/rotation of sunspots and interaction of filaments can destablize the magnetic field, and trigger the solar eruptive phenomena, {\\it e.g.}, flares, CMEs {\\it etc.} \\cite{min2009,kumar2010}.  The S-shaped {\\it or} inverted S-shaped sigmoids indicate the twisted field lines in the solar atmosphere.  The solar eruptions usually take place due to the increase of the twist in the magnetic field configurations of an active regions \\cite {canfield1999}. Recently, the MHD models of magnetic flux tubes show that the twist [$\\phi$] of \\,2.5--\\,3.5$\\pi$ is sufficient for solar eruptions, and the stable equilibrium of the magneto-fluid breaks when the total twist in the associated flux tubes crosses this critical value \\cite{fan2003,kliem2004,torok2004}.    \\begin{figure} \\centering \\includegraphics[width=8cm]{mdi_box3.pdf}  \\includegraphics[width=5.5cm]{sot_box1.pdf} \\includegraphics[width=5.1cm]{sot_box2_new.pdf} \\caption{Top: SOHO/MDI image of active region NOAA 10960 on 04 June 2007. The positive-polarity sunspot indicated by arrow plays an important role in triggering the M8.9/3B solar flare. The enlarged view of the sunspot group of the active region as indicated by a box in SOHO/MDI image, is shown in the SOT/blue continuum (4504 \\AA) image (bottom-left panel). A more closer view of the positive-polarity sunspot as indicated by a box in the SOT/blue continuum image, is also shown in the bottom-right image. Penumbral filaments, twisted in the counterclockwise direction, are clearly evident in this image.} \\label{fig1} \\end{figure} \\inlinecite{nandy2008} has shown the formation, evolution, and ejection of the magnetic flux ropes, which originate in the twisted magnetic structures through the combined action of surface flux transport processes (such as diffusion, meridional circulation and differential rotation). \\inlinecite{ishii1998} have investigated a flaring active region NOAA 5395 during March 1989 and found some peculiar vortex-like motion of small, satellite sunspots, which successively emerged from the leading edge of the sunspot group. They propose a schematic model of successive emergence of twisted and winding magnetic-flux loopdf coiling around a trunk of magnetic-flux tube, and concluded the flare triggering due to this emerging flux bundles.  The twisted flux-tube model \\cite{amari2000} and flux-injection driven model \\cite{krall2001}, both suggest that a twist-enhanced flux rope can play a crucial role in large-scale eruptive events. The primary mechanism for driving such eruptions may be due to the catastrophic loss of  MHD equilibrium \\cite{lin2000}. They have suggested that a flux rope is allowed to escape with a fairly small reconnection rate in the vertical current sheet created below. The highly twisted flux tubes store magnetic energy, which is necessary for the heating and particle acceleration during the solar eruptions. \\inlinecite{wang2002} have found a rapid disappearance of a sunspot associated with the M2.4 flare from NOAA AR 9830 on 20 February 2002 with hard X-ray sources located near this disappeared sunspot.  \\inlinecite{ishii2000} have also pointed out that the occurrence of high flare activity is restricted to the location and the time at which the strongly twisted magnetic flux ropes emerged to the photosphere in the long lived and large active regions, {\\it e.g.}, NOAA AR 4201.  The observational evidences of twisted helical structure are less abundant in the solar atmosphere. However, they may be an efficient mechanism for the triggering of solar eruptive phenomena  e.g. flares and coronal mass ejections (CMEs).  \\inlinecite{gary2004} and \\inlinecite{liu2003}  have observed the helical magnetic flux tubes with multiple turns, which were associated with double flares and CMEs. This observational evidence of the activation of helical flux tubes, and thus associated destabilization of large-scale magnetic fields of active region, may be important clues for the energy build-up processes of solar flares. However, we do not have sufficient understanding and observational signature of the twisted flux ropes and their evolution from the sub-photospheric level into the corona. There are a few observational signatures related to the generation of the twist in the solar filaments, which causes the disruption of their stable magnetic field configuration and generates solar eruptive events ({\\it e.g.}, \\opencite{liu2009}; \\opencite{williams2005} and references cited there). \\inlinecite{rust2005} have also found evidence of sigmoids in the solar corona which were governed and energized by the magnetic twist, without any large-scale destabilization of magnetic fields and associated eruptions from the Sun. \\inlinecite{gerrard2002} have found that the foot-point twisting motion may also generate the twisting in the active-region loopdf which trigger flare event after their reconnection with the surrounding opposite-polarity field lines.  AR 10960 shows the successive activation of helical twisted magnetic structures on 4 June 2007, which do not cause any eruption. Recently, evidence of the kink instability has been found in the right-handed twisted loop, which causes the B5.0 class flare in this AR during \\,04:40--04:51\\, UT \\cite{sri2010}. In the present paper, we study the M8.9/3B flare event of the same active region NOAA 10960 during \\,05:06--05:16\\, UT on 4 June 2007 using multi-wavelength observations. We find   rare observational evidence of the activation of helical-twisted magnetic structure in the active region, which may produce the M-class flare.  In Section 2, we present multi-wavelength observations of AR 10960 and the associated flare. In later sections, we present the discussion and conclusions. ",
        "conclusions": "We find multiwavelength evidence of the successive activation of helical twists that may help in the energy build-up process at the flare site. The built energy is released later in the form of M-class flare after secondary activation of this critical twist in the flux tube/rope and its reconnection with neighbouring opposite fields. The main conclusions of this study may be summarized as:  {\\it i)}. We report the dynamics of a single positive-polarity sunspot having twisted penumbral filament structure and successive activation of twisted helical magnetic structures.  {\\it ii)}. The activation of two helical structures/ropes played an important role in destabilizing the field lines and in triggering the flare. The twist in the secondary magnetic structures crosses the threshold limit (2.5\\,--\\,3.5$\\pi$), which probably produces the kink instability in this structure. The energy-release site of the M-class flare {\\it i.e.} at the center of the AR, coincides with the activated twisted magnetic structures. \\inlinecite{ishii1998} pointed out that the successive emergence of helical flux bundles plays crucial role in triggering flares and showed it by a schematic cartoon, but here we provide the observational evidence of the same.  {\\it iii)}. The M-class flare shows agreement with the quadrupolar (closed-closed) reconnection model (breakout) between two closed field lines \\cite{anti1998}. The asymmetric evolution is driven by footpoint shearing of one side arcade, where reconnection between the sheared arcade and the neighboring (unsheared) flux system triggers the flare. As the twisted magnetic structure moves away from the reconnection site, the flare intensity increases and reaches to maximum. This reveals the progressive reconnection of the twisted magnetic structure with the surrounding opposite-polarity fields.  {\\it iv)}. Penumbral disappearance during the decay phase and after of the flare event suggests that the magnetic-field structure to becomes more vertical. This indicates that the magnetic field changes from a highly inclined to an almost vertical configuration during the decay phase of the flare (just after the flare maximum), {\\it i.e.} part of penumbral magnetic field is converted into umbral fields. Our results are in agreement with previous studies \\cite{wang2004,liu2005}.  {\\it v)}. The decrease of the overlying field with height is a main factor in deciding whether the kink instability leads to a confined event or to a CME \\cite{torok2005}. In AR 10960, the slow variation of vertical component of the magnetic field with height is found to be the most likely the cause for the failed eruption during M-class flare.   \\begin{acks} We thank the reviewers for their valuable suggestions which improved the manuscript considerably. We acknowledge SOHO/MDI and TRACE for providing the data used in this study.  SOHO is a project of international cooperation between ESA and NASA. We acknowledge the Hinode and STEREO missions for providing us the high-resolution data. Hinode is a Japanese mission developed and launched by ISAS/JAXA, with NAOJ as domestic partner and NASA and STFC (UK) as international partners. It is operated by these agencies in co-operation with ESA and NSC (Norway). This work was supported in part by the Department of Science and Technology, Ministry of Science and Technology of India and in part by the Russian foundation for Basic Research (grants 09-02-00080 and 09-02-92626, INT/RFBR/P-38). \\end{acks}"
    },
    "1005/1005.2187_arXiv.txt": {
        "abstract": "{ Local luminosity functions are fundamental benchmarks for high-redshift galaxy formation and evolution studies as well as for models describing these processes. Determining the local luminosity function in the submillimeter range can help to better constrain in particular the bolometric luminosity density in the local Universe, and \\textit{Herschel} offers the first opportunity to do so in an unbiased way by imaging large sky areas at several submillimeter wavelengths. % We present the first \\textit{Herschel} measurement of the submillimeter $0<z<0.2$ local luminosity function and infrared bolometric (8-1000 $\\mu$m) local luminosity density based on SPIRE data from the HerMES \\textit{Herschel} Key Program over 14.7 deg$^2$. % Flux measurements in the three SPIRE channels at 250, 350 and 500 $\\mu$m are combined with \\textit{Spitzer} photometry and archival data. We fit the observed optical-to-submillimeter spectral energy distribution of SPIRE sources and use the $1/V_{max}$ estimator to provide the first constraints on the monochromatic  250, 350  and 500 $\\mu$m as well as on the infrared bolometric (8-1000 $\\mu$m) local luminosity function based on \\textit{Herschel} data. % We compare our results with modeling predictions and find a slightly more abundant local submillimeter population than predicted by a number of models. Our measurement of the infrared bolometric (8-1000 $\\mu$m) local luminosity function suggests a flat slope at low luminosity, and the inferred local luminosity density, $1.31_{-0.21}^{+0.24}\\,\\cdot\\,10^8\\,L_{\\odot}\\mathrm{Mpc^{-3}}$, is consistent with the range of values reported in recent literature. % } ",
        "introduction": "Local luminosity functions (LLFs) are essential for the interpretation of the evolution of cosmic sources at any wavelength. While the determination of high-redshift LFs requires deep observations, LLFs can exploit shallower wider-area sky maps. In addition, the extremely careful multi-wavelength identification work required for high-redshift sources is greatly simplified by the lower areal density of bright local sources. The pioneering exploration of the IR sky by the IRAS satellite revealed the rapid evolution of the infrared (IR) LLF (Saunders et al. 1990), illustrating the importance of local studies at infrared (IR) wavelengths. ISO was later able to follow the evolution of the IR LF up to $z\\simeq1$ (Pozzi et al. 2004), while \\textit{Spitzer} studies using the MIPS 24 $\\mu$m (Le Floc'h et al. 2005, Marleau et al. 2007, Rodighiero et al. 2010) and 70 $\\mu$m (Frayer et al. 2006, Huynh et al. 2007, Magnelli et al. 2009) channels have both extended our reach to higher redshifts and improved our constraints on the IR bolometric (8-1000 $\\mu$m) luminosity of detected sources. Due to the moderate volume per unit sky area that one can probe at low redshifts, an essential pre-requisite for determining the LLF is the imaging of large fields. This has so far been difficult for ground-based submillimeter (SMM) cameras, and for example Dunne et al. (2000) carried out SCUBA observations of IRAS bright galaxies instead. It was only recently that BLAST was able to map large areas at SMM wavelengths and thus to first determine the evolution of the SMM LLF (Eales et al. 2009). The much improved mapping speed and angular resolution of the SPIRE (Griffin et al. 2010) instrument onboard the \\textit{Herschel} Space Observatory (Pilbratt et al. 2010) is exploited by the \\textit{Herschel} Multi-tiered Extragalactic Survey (HerMES, \\url{http://hermes.sussex.ac.uk}) Key Program (Oliver et al. 2010, in prep), which will eventually cover 70 deg$^2$. While other early \\textit{Herschel} studies are probing the evolution of SMM number counts and LFs up to high redshifts (Oliver et al. 2010, Eales et al. 2010), we present a first assessment of the LLF as measured in the three SPIRE bands at 250, 350 and 500 $\\mu$m, providing the first robust constraints on the Local Luminosity Density (LLD) at these wavelengths. This study will thus establish a reliable local benchmark at a very early stage of the \\textit{Herschel} mission against which models and high-redshift studies can be compared. Throughout, we adopt a standard cosmology with $\\Omega_M=0.28$, $\\Omega_\\Lambda=0.72$, $H_0=72~\\mathrm{km\\,s^{-1}\\,Mpc^{-1}}$. ",
        "conclusions": "Using SPIRE imaging from the HerMES Key Program covering 14.7 deg$^2$ (about 20\\% of the total SPIRE area eventually to be covered by HerMES) we presented the first \\textit{Herschel} measurement of the SMM LLF. We computed the IR monochromatic 250, 350 and 500 $\\mu$m as well as the IR bolometric (8-1000 $\\mu$m) LFs of SPIRE $0<z<0.2$ sources and compared them with modeling predictions for IR LLFs. Relying on \\textit{Spitzer} 24 $\\mu$m positions to identify SMM sources and on optical spectroscopic surveys to determine their distances, we placed solid statistical constraints on the SMM LLF and LLD at a very early stage of the \\textit{Herschel} mission and compared our results with pre-\\textit{Herschel} predictions. We found a slightly more abundant local submillimeter population than predicted by a number of models, a flat slope of the IR LLF at low luminosity and an IR bolometric LLD of $1.31_{-0.21}^{+0.24}\\,\\cdot\\,10^8\\,L_{\\odot}\\mathrm{Mpc^{-3}}$, toward the higher end of values from recent literature. Once HerMES observations are completed, we will be able to better probe the faint and bright end of the SMM LLF and study its evolution with redshift in much greater detail."
    },
    "1005/1005.2536_arXiv.txt": {
        "abstract": "As planetary embryos grow, gravitational stirring of planetesimals by embryos strongly enhances random velocities of planetesimals and makes collisions between planetesimals destructive. The resulting fragments are ground down by successive collisions.  Eventually the smallest fragments are removed by the inward drift due to gas drag.  Therefore, the collisional disruption depletes the planetesimal disk and inhibits embryo growth.  We provide analytical formulae for the final masses of planetary embryos, taking into account planetesimal depletion due to collisional disruption.  Furthermore, we perform the statistical simulations for embryo growth (which excellently reproduce results of direct $N$-body simulations if disruption is neglected).  These analytical formulae are consistent with the outcome of our statistical simulations.  Our results indicate that the final embryo mass at several AU in the minimum-mass solar nebula can reach about $\\sim 0.1$ Earth mass within $10^7$ years.  This brings another difficulty in formation of gas giant planets, which requires cores with $\\sim 10$ Earth masses for gas accretion. However, if the nebular disk is 10 times more massive than the minimum-mass solar nebula and the initial planetesimal size is larger than 100 km, as suggested by some models of planetesimal formation, the final embryo mass reaches about 10 Earth masses at 3--4\\,AU.  The enhancement of embryos' collisional cross sections by their atmosphere could further increase their final mass to form gas giant planets at $5$--$10$\\,AU in the solar system. ",
        "introduction": "In the standard scenario of planetary formation, terrestrial planets and cores of Jovian planets are formed through the accretion of planetesimals with initial size of 10--100\\,km (e.g., Hayashi et al. 1985).  This process called planetary accretion has been investigated by statistical simulations (e.g., Wetherill and Stewart 1993; Inaba et al. 2003; Kenyon and Bromley 2004), by $N$-body simulations (e.g., Kokubo and Ida 1996; 2002), and by the hybrid method that combines an $N$-body simulation for large bodies called planetary embryos with a statistical simulation for small bodies (e.g., Bromley and Kenyon 2006; Chambers 2008).  Assuming that each collision of bodies leads to a perfect agglomeration, the growth of planetesimals can be accurately computed with $N$-body codes.  Inaba et al. (2001) showed that the evolution of mass distribution and velocity dispersion of bodies calculated by $N$-body simulation were reproduced by the statistical simulation applying the collision rate and the velocity-dispersion-evolution rate based on the results of orbital integrations. Therefore, statistical simulations are reliable unless planetary embryos collide with each other after developing a long-term orbital instability.  Planetary embryos initially form through the runaway growth (e.g., Wetherill and Stewart 1989; Kokubo and Ida 1996).  The embryos keep their orbital separations and grow through collisions with surrounding planetesimals (Kokubo and Ida 1998).  At this stage, planetesimals of almost initial size dominate the total mass of bodies (surface density). When embryos reach about the mass of Mars, the velocity dispersion of planetesimals is increased by the gravitational scattering at the embryos.  Accordingly, another growth regime sets in, referred to as oligarchic growth.  At this stage, higher velocities of planetesimals cause their collisional fragmentation.  After a chain of successive destructive collisions, often called ``collision cascade'', bodies get smaller and smaller until they are removed by gas drag in protoplanetary disks or by radiation pressure and/or Poynting-Robertson drag in debris disks.  As a result, collision cascade decreases the surface density, which slows down the growth of planetary embryos (Inaba et al. 2003; Kenyon and Bromley 2008).  The planetary core (embryo) exceeding the critical core mass as large as 10 Earth masses can no longer retain a hydrostatic envelope, resulting in the gas accretion and formation of the gas giant planets (e.g., Mizuno et al.\\ 1980; Bodenheimer \\& Pollack 1986; Pollack et al.\\ 1996; Ikoma et al.\\ 2001).  However, since the embryo growth is hindered by the planetesimal depletion in collision cascade, the final embryo can hardly grow beyond the Mars mass and thus cannot form a gas giant (Kobayashi and Tanaka 2010).  However, it is possible that fragmentation inhibits the embryo growth to a lesser extent than ascertained before.  Collisional fragments orbiting a central star drift inward by gas drag.  The drift time shortens as bodies become smaller by collision cascade, until at a certain size their motion gets coupled with gas motion.  Such coupled bodies have lower drift velocity and thus can survive around planetary embryos for longer time.  Kenyon and Bromley (2009) suggested that accretion of those coupled bodies by planetary embryos may promote further growth of the latter.  In their simulations, they assumed that the coupled bodies no longer experience an inward drift.  However, the drift, although at a reduced rate, is important to determine the embryo mass gain due to the accretion of coupled bodies.  This paper investigates the embryo growth taking into account the accretion of fragments resulting from the collision cascade before their removal by gas drag.  To this end, we perform an analytic study that extends the model of Kobayashi and Tanaka (2010) by including the removal by gas drag, as well as statistical simulations.  In the numerical treatment, we do not neglect the drift of coupled bodies by gas drag.  Instead, we take into account that the gas drag law changes for small bodies in the both analytical and numerical procedure (e.g., Adachi et al. 1976).  The goal is to find out to what extent the collisional fragmentation combined with gas drag would affect the embryo growth and whether an embryo can reach the critical core mass.  We develop the analytic theory and derive the final embryo mass with fragmentation and gas drag in Section 2.  In Section 3, we check the formulae for the final mass against the statistical numerical simulations. Section 4 contains a summary and a discussion of our results. ",
        "conclusions": "In this paper, we studied analytically and numerically the growth of planetary embryos in the framework of the standard scenario. We took into account that embryos growing in the oligarchic mode pump up relative velocities of planetesimals, which causes their collisional fragmentation. We also considered the fate of smallest fragments of the resulting collision cascade, namely their inward drag in the ambient gas and possible accretion by nascent embryos.  Taking into account fragmentation, we have analytically derived the final mass of a planetary embryo in Section \\ref{sec:theo} in two cases: the mass $M_{\\rm c}$ of an embryo growing through the accretion of planetesimals which are removed by collision cascade and the mass $M_{\\rm f}$ in the case where an embryo grows through collisions with fragments which are removed by gas drag.  In Sec.~\\ref{eq:num}, we showed that the final embryo mass obtained in numerical simulations can be reproduced by the larger of the two analytic estimates, MAX($M_{\\rm c}, M_{\\rm f}$).  If the embryo mass $M_{\\rm n}$ of Eq.~(\\ref{eq:m_nogas}) for the case without fragmentation is smaller than MAX($M_{\\rm c}, M_{\\rm f}$), fragmentation is negligible and the final mass is given by $M_{\\rm n}$.  Altogether, our analytical formulae provide a good estimate of the final mass of planetary embryos.  The gas giant planet formation via core accretion requires the solid core as large as $\\sim 10 M_{\\oplus}$ (e.g., Ikoma et al.~2001).  Kenyon and Bromley (2009) suggested that the planetary embryo may further grow by the accretion of bodies coupled with gas because the coupled bodies are no longer involved in collision cascade.  Taking into account three regimes of gas drag (see Eq.~\\ref{eq:cd_adopt}), we found that the collision cascade halts at bodies having $v^2 \\sim Q_{\\rm D}^*$, i.e. at sizes that are larger than the critical coupling size corresponding to $\\tilde{\\tau}_{\\rm stop} = 1$.  As a result, coupled bodies should be produced only in little amounts and can hardly contribute to planetary growth, despite their low velocity dispersion in the laminar disks considered here.  In this paper, we derive the final embryo masses, assuming a uniform distribution of solid bodies.  However, the planetary embryos can open gaps in a solid disk, which would affect the accretion of fragments on the embryo and their removal by the gas drag.  Tanaka and Ida (1999) derived a condition for the gap opening around a migrating embryo. Replacing the embryo's migration time by the drift time $\\tau/2\\eta^2$ for fragments, we can obtain a criterion for an embryo with mass $M$ to open a gap in the solid disk from their Eq.~(3.6), \\begin{equation} \\frac{\\tau \\Omega_{\\rm K} h_M^2}{4 \\pi \\eta^2} \\geq 0.81 \\left[\\sqrt{1+0.45(\\tau\\Omega_{\\rm K}/2\\pi)^{2/3}} +1 \\right]^2,\\label{eq:gap_crit} \\end{equation} where $\\tau$ is given by Eq.~(\\ref{eq:tau_gas}) for fragments surrounding the embryos.  Since Eq.~(\\ref{eq:gap_crit}) requires $\\tau\\Omega_{\\rm K}/2\\pi \\ll 1$ for $M \\ga 0.1 M_\\oplus$, the critical fragment radius $r_{\\rm g}$ above which Eq.~(\\ref{eq:gap_crit}) is satisfied is given by \\begin{equation} r_{\\rm g} \\simeq 1.0 \\times 10^2 \\left(\\frac{M}{0.1 M_{\\oplus}}\\right)^{-1/3} \\Biggl(\\frac{a}{5\\,{\\rm AU}}\\Biggr)^{3/8}\\, \\rm{m}, \\label{eq:rg} \\end{equation} where we adopt $q=3/2$ and $C_{\\rm D} = C_{\\rm C, Stokes}$.  Levison et al.~(2010) performed $N$-body simulations involving fragments and found $r_{\\rm g} \\sim 30$\\,m for $M \\geq M_{\\oplus}$, which is consistent with Eq.~(\\ref{eq:rg}). On the other hand, as we have shown, the solid surface density is reduced by the elimination of fragments at the low-mass end of collision cascade due to gas drag.  In addition, the embryo grows through the accretion of the fragments for small initial planetesimals.  The radius $r_{\\rm e}$ of the fragments is estimated by Eq.~(\\ref{eq:tstop_est}) as \\begin{equation} r_{\\rm e} \\simeq 5.0 \\left(\\frac{M}{0.1 M_{\\oplus}}\\right)^{-1/2} \\Biggl(\\frac{a}{5\\,{\\rm AU}}\\Biggr)^{1/2}\\, \\rm{m}.\\label{eq:re} \\end{equation} Eqs.~(\\ref{eq:rg}) and (\\ref{eq:re}) give $r_{\\rm e} \\ll r_{\\rm g}$. This means that the collision cascade quickly grinds planetesimals down to the size much less than $r_{\\rm g}$. As a result, the gap formation in the solid disk composed of fragments with the radius $r_{\\rm e}$ does not occur and the embryo growth would not be influenced.  Considering the range of initial planetesimal radii $r_0 = 1$--100\\,km, and the range of disk surface densities $\\Sigma_{1} \\sim 7$--$70\\,{\\rm g/cm}^2$ at 3.2\\,AU (1--$10\\times$MMSN), our results suggest that the fragmentation averts the planetary growth at $M \\sim 0.1$--$10\\,M_\\oplus$ at several AU.  This is consistent with the results of planetary growth simulations beyond 30\\,AU by Kenyon and Bromley (2008).  The cores approaching the critical mass of $10\\, M_\\oplus$ can only form at $a=3$--4\\,AU, only from planetesimals that have $r_0 \\sim 100\\,$km initially, and only in densest nebula with $\\Sigma_{1} \\sim 70\\,{\\rm g/cm}^2$.  Thus it appears problematic to explain formation of giant planets at 5--10\\,AU, such as Jupiter and Saturn in the solar system, even for $\\Sigma_{1} \\sim 70\\,{\\rm g/cm}^2$ and $r_0 \\sim 100$\\,km.  An effect that could help further is the enhancement of the embryo's collisional cross section due to a tenuous atmosphere acquired by the embryo in the ambient gas.  Inaba and Ikoma (2003) and Inaba et al. (2003) have shown that a core with mass exceeding about $0.1 M_\\oplus$ could grow by that effect to the critical core mass.  As found here, the embryo masses can indeed be about $0.1 M_\\oplus$ under the minimum mass solar nebula conditions, at least for an initial planetesimal radius greater than $10$\\,km.  On any account, our results strongly suggest that increasing the disk surface density and/or the initial embryo size helps forming larger cores.  This brings up the question whether the values that we need to grow the embryo to the critical mass, $10\\times$MMSN and $r_0 \\sim 100$\\,km~--- or less, if the atmospheric effect is taken into account~--- appear plausible.  As far as the disk density is concerned, the answer is probably yes.  $10\\times$MMSN is close to the gravitational stability limit, and all densities below this limit cannot be ruled out. For instance, Desch (2007) and Crida (2009) point out that the MMSN profile is inconsistent with the ``Nice model'' (Gomes et al. 2004) and should be replaced with another surface density profile, which would possibly imply larger surface densities at several AU.  We now consider the initial size of planetesimals. The mechanisms of planetesimal formation are highly debated but, despite intensive effort, remain fairly unknown. Classical models in which dust smoothly grows to planetesimals with $r_0 \\sim 1$\\, km face the ``meter barrier'' problems: first, meter-sized objects should be lost to the central star as a result of gas drag (Weidenschilling 1977, Brauer et al. 2008), and second, further agglomeration of meter-sized objects upon collision is problematic (Blum and Wurm 2008).  Accordingly, in the last years, competing scenarios were suggested that circumvent these barriers, such that the ``primary accretion'' mechanism proposed by Cuzzi et al. (2008) and ``graviturbulent'' formation triggered by transient zones of high pressure (Johansen et al. 2006) or by streaming instabilities (Johansen et al. 2007).  These models all imply rapid formation of rather large planetesimals in the $r_0 \\sim 100$--$1000$\\, km range.  Support for these scenarios may come from the analysis of left-over planetesimals in the solar system.  Morbidelli et al. (2009), for instance, suggest that the initial planetesimals should be larger than 100\\,km to reproduce the mass distribution of asteroids in the main belt.  Our last remark is related to the Type I migration of bodies due to interaction with gas (e.g., Tanaka et al. 2002).  The planetary embryo grows in the runaway mode followed by the oligarchic one.  The timescale of the runaway growth is proportional to the initial size of bodies. In addition, the embryos rapidly grow through collisions with small fragments for small $r_0$ even in the the oligarchic growth.  The embryo eventually grows to a larger mass if it is large initially, while it forms earlier if it starts with a smaller mass.  Therefore, initially small planetesimals seem to be preferred for the core formation prior to their removal due to Type I migration.  However, this is only valid for laminar disks.  Turbulence may help prevent the loss of growing embryos to Type I migration (e.g., Laughlin et al. 2004).  So, the alternative scanarios of planetesimal formation that all require turbulent disks may help here, too."
    },
    "1005/1005.2193_arXiv.txt": {
        "abstract": "It has recently been suggested that conduction-driven magnetohydrodynamic (MHD) instabilities may operate at all radii within an intracluster medium (ICM), and profoundly affect the structure of a cluster's magnetic field.  Where MHD instabilities dominate the dynamics of an ICM, they will re-orient magnetic field lines {\\it perpendicular} to the temperature gradient inside a cooling core, or {\\it parallel} to the temperature gradient outside it.  This characteristic structure of magnetic field could be probed by measurements of polarized radio emission from background sources. Motivated by this possibility we have constructed 3-d models of a magnetized cooling core cluster and calculated Faraday rotation measure (RM) maps in the plane of the sky under realistic observing conditions.  We compare a scenario in which magnetic field geometry is characterized by conduction driven MHD instabilities to that where it is determined by isotropic turbulent motions. We find that future high-sensitivity spectro-polarimetric measurements of RM, such as will be enabled by the {\\it Expanded Very large Array} and {\\it Square Kilometer Array} can distinguish between these two cases with plausible exposure times. Such observations will test the existence of conduction-driven MHD instabilities in dynamically relaxed cooling core clusters. More generally, our findings imply that observations of Faraday RM should be able to discern physical mechanisms that result in qualitatively different magnetic field topologies, without a priori knowledge about the nature of the processes. ",
        "introduction": "Understanding the role and evolution of magnetic fields in clusters of galaxies is of significant importance for many questions including the origin of cluster magnetic fields, the interaction of active galactic nuclei (AGN) with the intracluster medium (ICM), and physical processes operating within the ICM plasma.  One of the key techniques used to obtain information about the strength and structure of cluster magnetic fields is the analysis of Faraday rotation from polarized radio sources located behind and within clusters. The Faraday effect rotates the polarization plane of electromagnetic waves traveling through a magnetized plasma, by an amount proportional to the (square of the) wavelength, plasma density, and the strength of the field in the direction of propagation.  The sources' intrinsic polarization need not be known, as the effect can be observed as a characteristic wavelength-dependent rotation measure (RM) signature. Observations of a few nearby clusters have established the presence of magnetic fields with typical strengths of ${\\rm few\\,\\mu G}$ in non-cool core clusters and in excess of 10${\\rm\\, \\mu G}$ in the centers of cool core clusters \\citep{ct02,gf04,fg08}. Detailed high resolution RM images of radio galaxies in merging and cooling core clusters indicate that the RM distribution is characterized by patchy structures of a few kpc in size. Furthermore, the RM distribution appears in general in good agreement with the RM expected for magnetic fields geometry characterized by turbulent motions with a power-law power spectrum \\citep[][for e.g.]{ev03b,murgia04, guidetti08,laing08,bonafede10}. This base of knowledge will be greatly enhanced by radio observatories such as the {\\it Expanded Very Large Array}\\footnote{http://science.nrao.edu/evla} ({\\it EVLA}) and {\\it Square Kilometer Array}\\footnote{http://www.skatelescope.org} ({\\it SKA}), which will provide the sensitivity to study both background and embedded sources with sufficient density to map out the RM patterns in many and more distant clusters. This will allow to study the structure of magnetic field in ``quiescent'' regions of clusters, away from radio galaxies, as well as the effect of the radio-mode feedback in their vicinity.  The interpretation of these results can strongly benefit from advances in the theoretical understanding of the dynamics of weakly magnetized, low density plasmas. The ICM is a dilute plasma, i.e., the gyro-frequency of both the electrons and the ions is much greater than the collision frequency. The principal effect of this is highly anisotropic thermal conduction that fundamentally alters the classical Schwarzschild criterion for convection \\citep{balbus00}. Under these circumstances, the outer regions of clusters where temperature decreases with radius may become susceptible to a conduction-driven {\\it magneto-thermal instability} (MTI). The properties of this instability have been studied in MHD simulations \\citep{ps05,ps07,psl08}. \\citet{psl08} found that MTI can profoundly affect the temperature distribution in the outer regions of a cluster in only a few billion years, by re-orienting the lines of magnetic field to be preferentially radial, resulting in a thermal conduction at a high fraction of the Spitzer conductivity.  The temperature profile in the centers of cooling-core galaxy clusters makes them stable to MTI. However, \\citet{quataert08} found that cooling cores are characterized by a sister {\\it heat-flux buoyancy instability} (HBI) that arises in regions where temperature increases with radius. An initial investigation of nonlinear evolution of the HBI \\citep{pq08} indicates that its primary effect is to re-orient the lines of magnetic field perpendicular to the temperature gradient, and thus strongly suppress heat conduction. The properties of HBI have been investigated in global 3-d models of isolated cooling core clusters \\citep{pqs09,bogdanovic09}. They suggest that, once the magnetic field lines have been wrapped into spherical surfaces surrounding the core, the effective thermal conduction is suppressed to a small fraction of the Spitzer value, leading to insulation of the core from further conductive heating, and to a subsequent thermal collapse.  Very recently, \\citet{br10} discovered an associated pair of {\\it overstabilities} that affect precisely those configurations that are stable to the well-established HBI and MTI. They predict that configurations which tend to result from the non-linear evolution of the HBI have $g$-modes that are driven overstable by radiative loses. On the other hand, configurations which tend to result from the non-linear evolution of the MTI have $g$-modes that are driven overstable by the conductive heat flux. The effects of these overstabilities for the ICM plasma thermodynamics and the properties of magnetic field are yet to be understood. We do not consider them in this work and instead focus on MTI and HBI instabilities.  The combination of MTI and HBI instabilities should lead to a characteristic structure of magnetic fields in some cooling core clusters.  If these conduction-driven instabilities dominate the dynamics of the ICM, the magnetic field lines would be preferentially oriented radially in the outer region and azimuthally within the cooling core. This may not be the case in all clusters -- it has recently been shown that MHD instabilities can be overwhelmed by even moderate levels of driven turbulence \\citep{ro09,pqs10}, such as might result from sub-cluster mergers, motions of cluster member galaxies, or various forms of AGN feedback. Which (if any) physical mechanism dominates in the ICM depends sensitively upon the magnitude and distribution of turbulence, which is currently only poorly understood. Motivated by this question and guided by the MHD simulations, we have constructed 3-d models of clusters for two distinct cases and simulated the RM maps that might be observed with high sensitivity radio polarization measurements.  This paper is organized as follows. We describe the ingredients of our cluster models in \\S~\\ref{S_model} and the main properties of derived RM maps in \\S~\\ref{S_RMmaps}. We discuss the importance of depolarization and diffuse emission in \\S~\\ref{S_depolarization}, then present the discussion with conclusions in \\S~\\ref{S_discussion}. ",
        "conclusions": "\\label{S_discussion}  Even if MTI and HBI instabilities operate uninhibited in real clusters, the radial-azimuthal field geometry is likely to be perturbed by intermittent phases of AGN activity and mergers. As these perturbations are attenuated over time, they can source a new cycle of MHD instabilities. The relics of such events, including shocks, bubbles, and ridges of magnetic field lines swept by intracluster galaxies, may also be recognizable in high quality RM maps -- testifying to the past evolution of a cluster. But with a characteristic time scale for the saturation of MHD instabilities of only a few billion years, the magnetic field lines can be driven to recover their orientation relatively quickly. As a consequence, a fraction of clusters may exhibit preferential magnetic field geometries, despite episodic disruptions.  Circumstantial evidence in support of this hypothesis may already exist in several forms.  Recent detailed observations of the Virgo cluster revealed an isothermal region with remarkable azimuthal symmetry cocooned in the cool core of the cluster between the active radio-lobes of M87 \\citep{million10}.  The metallicity of the isothermal gas in the same region is non-uniform and clumpy. It is reasonable to expect that the high metallicity parcels of gas were uplifted from the low entropy, cool core region of the cluster, where the density of stars is highest. That the parcels now have the entropy and temperature of the surrounding ambient plasma, but retain clumpy metallicity distribution, suggests that turbulent mixing was not efficient here and that the gas was instead heated by conduction.  The presence and shape of the isothermal front indicate that the magnetic field geometry in this region may be {\\it predominantly azimuthal}. Furthermore, the absence of efficient turbulence indicates that HBI may operate unhindered in this region and that the instability may be more robust than suggested by some recent theoretical works \\citep{ro09, pqs10}. However, at least one more physical phenomenon may produce a similar magnetic field topology: \\citet{ro09} find that weak turbulent motions lead to trapped $g$-modes and result in gas motions that are preferentially tangential. They suggest that magnetic fields in clusters experiencing such $g$-modes can in principle become tangential even in the absence of thermal conduction and the HBI. While some thermal conduction seems to be implied in the case of the Virgo cluster, the concurrent presence of $g$-modes cannot be eliminated.  \\citet{juett10} draw attention to what may be more circumstantial evidence for MHD instabilities.  As much as $20\\%$ of the sample of 70 clusters presented by \\citet{snowden08} have puzzling temperature profiles that appear to be quasi-isothermal at $\\sim$Mpc radii\\footnote{The classification of the sample of quasi-isothermal clusters according to the thermal state of their cores has not been reported thus far.}. The clusters show no signs of interaction, which could possibly have accounted for their temperature distribution, and they appear dynamically relaxed. Moreover, it is unclear whether models based on theoretical studies of cluster properties could account for this unusual class of objects \\citep{nagai07}. If the class is shown to be unexpected, it may point to some missing physics in cosmological simulations. A possible explanation for the temperature structure of these objects is that thermal conduction operates very efficiently in their outer regions, an effect that arises as a natural consequence of the MTI instability. Radial configuration of magnetic field lines in outer regions of such clusters could only persist if they are isolated and unperturbed for sufficiently long periods of time, consistent with the observed properties of the quasi-isothermal clusters. Another line of evidence comes from the study of magnetic field structure around galaxies in the Virgo cluster. \\citet{pd09} find that in Virgo, which seems to be in transition to a cool core, the global magnetic field has a predominantly radial orientation at large radii, which again suggests the operation of MTI in its ICM.  In this study, we have evaluated the effects of two different physical mechanisms on the Faraday rotation measure of a magnetized cooling core cluster, in the context of the planned capabilities of the {\\it EVLA} and {\\it SKA} radio observatories. We compare a theoretical scenario in which conduction-driven MHD instabilities dominate the dynamics of the ICM, to a scenario in which magnetic field topology is defined by turbulent motions. We employ the polar shapelets image analysis method to efficiently detect patterns in the RM image with specific rotational symmetries, and thus classify their morphologies. Within the bounds of our simple models we find that the two mechanisms can produce strikingly different RM patterns and that future spectro-polarimetric measurements will have sufficient sensitivity to discriminate between them that can be achieved in longer exposures with {\\it EVLA} and relatively short exposures with {\\it SKA}. We propose that the effect of the HBI and MTI instabilities be sought for in dynamically relaxed cooling core clusters, and especially in the subclass of clusters with quasi-isothermal temperature profiles at large radii.  More generally, it should be possible to discern physical mechanisms that result in qualitatively different magnetic field topologies from observed Faraday rotation measure maps, without a priori knowledge about the nature of the processes. Such observations will enable detailed investigations into the behavior of MHD instabilities and other associated physical phenomena, which are of far reaching importance to a number of fundamental questions related to energy transport in clusters."
    },
    "1005/1005.5299_arXiv.txt": {
        "abstract": "We report the discovery of burst oscillations at the spin frequency in ten thermonuclear bursts from the accreting millisecond X-ray pulsar (AMXP) IGR~J17511--3057. The burst oscillation properties are, like those from the persistent AMXPs SAX J1808.4--3658 and XTE J1814--338, anomalous compared to burst oscillations from intermittent pulsars or non-pulsing LMXBs. Like SAX J1808.4--3658 they show frequency drifts in the rising phase rather than the tail.  There is also evidence for harmonic content. Where IGR J17511--3057 is unusual compared to the other two persistent pulsars is that oscillations are not detected throughout all bursts. As accretion rate drops the bursts get brighter and their rise/decay time scales become shorter, while the oscillation amplitude falls below the detection threshold: first in the burst peak and then also in the rise.  None of the bursts from IGR J17511--3057 show evidence for photospheric radius expansion (which might be expected to suppress oscillation amplitude) which allow us to set an upper limit to the distance of 6.9 kpc.  We discuss the implications of our results for models of the burst oscillation mechanism. ",
        "introduction": "\\label{sec:intro}   It is thought that many low-mass X-ray binaries (LMXBs) contain rapidly rotating neutron stars. However measuring the neutron star spin has been a challenge for more than 30 years. So far we have been able to measure the spin of only 25 out of more than 150 LMXBs \\citep[see, e.g.,][for a catalog]{Liu07} thanks to episodes of X-ray pulsations at the neutron star spin frequency. These episodes can last for periods as short as a few seconds or as long as days to weeks.  When pulsations occur in the persistent X-ray emission, the sources are known as accreting millisecond X-ray pulsars (AMXPs). Today there are 13 of these systems \\citep[see][for the latest discovery]{Altamirano10f,Markwardt10a} and their pulsations are thought to be powered by accretion, in a scenario in which matter from the accretion disk is channeled by the magnetic field lines onto the magnetic pole, forming a hot spot visible in X-rays.  When pulsations occur during surface thermonuclear explosions \\citep[known as Type I X-ray bursts, see, e.g][ for a review]{Strohmayer06}, the sources are known as nuclear-powered X-ray pulsars (NPXPs) or burst oscillation sources. Prior to the discovery of IGR J17511--3057 there were 14 confirmed NPXPs \\citep[including the four which are also AMXPs, and excluding those which require confirmation, see][]{Watts08,Altamirano10c,Galloway10}, of which two are persistent AMXPs and two are intermittent AMXPs (see below).  For clarity, in the rest of this work we will refer to the accretion-powered pulsations as ``pulsations'' and to the nuclear-powered pulsations as ``burst oscillations''.    Burst oscillations presumably develop when the burning or its aftermath becomes in some way asymmetric.  However the exact mechanism that generates burst oscillations is not yet known, and they are not observed in every burst even in sources that do show them. All theoretical models proposed to date fail to explain certain aspects of the data, and it seems increasingly likely that the magnetic field in the surface layers (largely neglected by current models) may play an important role.  In this regard systems that are both AMXPs and NPXPs are particularly useful.  The magnetic field in these sources is obviously dynamically important, since it channels accretion flow onto the stellar surface. This allows us to gauge field strength and any effect that channeling may have in establishing either temperature or fuel distribution gradients in the surface layers \\citep{Strohmayer06, Watts08a, Watts08b, Lamb09a}.   Prior to the discovery of IGR J17511--3057, there were only four systems that were both AMXPs and NPXPs.  If one compares the duty cycle of their pulsations, these systems can be divided into two separate groups. The first group contains the AMXPs SAX~J1808.4--3658 \\citep{Wijnands98} and XTE~J1814--338 \\citep{Markwardt03d, Strohmayer03a}; these systems show pulsations persistently during all their outbursts \\citep[see][respectively]{Hartman09,Papitto07}. The second group contains the AMXPs HETE~J1900.1--2455 \\citep{Kaaret06} and Aql X-1 \\citep{Casella08}; pulsations from these sources were only detected intermittently. HETE~J1900.1--2455 has been in outburst for almost 4.5 years, but showed pulsations only intermittently during the first few months of activity \\citep{Galloway07a,Galloway08a}.  The intermittent AMXP Aql X-1 is particularly extreme: pulsations were detected only for $\\approx150$~sec out of the $\\approx1.5$~Msec of available data \\citep{Casella08}.   These two groups of sources differ not only in their pulsation duty cycle, but also in the characteristics of their burst oscillations. The intermittent pulsars show burst oscillations with similar characteristics to those observed in all non-AMXPs \\citep{Zhang98b,Galloway08,Watts09}.  Burst oscillations are detected in only a subset of bursts.  They show no harmonic content and are detected more often during decay than burst rise.  Frequency drifts, which are common, are slow.  Burst oscillation amplitudes tend to increase with energy. The persistent pulsars, by contrast, are both rather anomalous.  They show burst oscillations in all bursts, throughout the whole burst (i.e. starting at the rise, continuing at the peak -- except during photospheric-radius expansion peaks (PRE)-- and during the decay). SAX J1808.4-3658 shows very large fast frequency drifts in the burst rise that overshoot the spin frequency \\citep{Chakrabarty03}, while XTE J1814-338 is remarkable for its frequency stability and phase-locking \\citep{Strohmayer03a, Watts05,Watts08a}.  XTE J1814-338 has burst oscillations whose amplitude rises with energy \\citep{Watts06a}, and the oscillations also have substantial harmonic content.  These differences support the idea that magnetic field plays an important part in determining burst oscillation properties, either directly or by generating composition/temperature/fuel depth gradients in the surface layers.  Understanding how burst oscillation properties vary between and within the groups of persistent AMXPs, intermittent AMXPs and non-AMXPs is the first step to shedding light onto the mechanism (or mechanisms) that produce brightness asymmetries during X-ray bursts.  In this paper we report the discovery of burst oscillations from the persistent AMXP IGR J17511--3057.  We present a detailed analysis of the burst oscillation properties and discuss how they fit into the existing population.   \\subsection{IGR~J17511--3057}\\label{sec:source}  IGR~J17511--3057 was first unambiguously identified by the INTEGRAL gamma ray mission on September 12th, 2009 \\citep{Baldovin09}. However, the source was initially detected one day before by RXTE PCA Bulge scans as a rising flux attributed to two known sources in the same 1$^\\circ$ field of view: XTE~J1751--305 and GRS~1747--312 \\citep[the former is a known $\\approx435$~Hz AMXP, see][]{Markwardt02}. The first PCA pointed observation was performed on September 12th, and revealed coherent pulsations at $\\approx$245 Hz \\citep{Markwardt09}. Quasi-simultaneously with the PCA observation, Swift observed in the direction of the AMXP XTE~J1751--305 and did not detect any source at its known Chandra position, or any other position in the 22 arcmin diameter FoV \\citep{Markwardt09}. All these observations together confirmed that the source in outburst was not the known AMXP XTE~J1751--305 but a new one. A preliminary orbital solution yielded a pulse frequency of 244.8337(1) Hz, an orbital period of 207.4(8) minutes, and a projected semi-major axis of 274(1) lt-ms.  The mass function was 0.00107(2) $M_{\\sun}$, giving a minimum companion mass of $\\approx0.13$ $M_{\\sun}$.  Swift observations at the beginning of the outburst detected thermonuclear (Type I) X-ray bursts \\citep{Bozzo09}. RXTE follow-up observations led to the detection of burst oscillations during these thermonuclear bursts \\citep{Watts09a}. Chandra-HETG observations were performed on September 22nd \\citep{Nowak09}, leading to a refined position ($\\alpha=17^h 51^m 08^s.66$, $\\delta = -30^o 57^{'} 41^{''}$). Using the new coordinates, \\citet{Riggio09} and \\citet{Papitto09a} reported an improved orbital ephemeris (using RXTE and XMM-Newton data, respectively). Very recently, \\citet{Papitto10} presented a detailed analysis of the X-ray spectra of IGR~J17511--3057 using 70~ksec of XMM-Newton data. Near infrared follow-up observations identified the counterpart at a magnitude of K$_s=18.0\\pm0.1$ \\citep{Torres09a,Torres09b}. Very recently, \\citet{Bozzo10} reported on the Swift monitoring observations during the outburst. ",
        "conclusions": "The burst oscillations of IGR~J17511--3057 are, like those of the other two persistent AMXPs with burst oscillations, atypical compared to burst oscillations from the intermittent and non-pulsars. Particularly notable are the frequency drifts and the fact that several of the burst oscillation trains have detectable harmonic content.  These differences between the AMXPs and the other sources suggest that the presence of a channeling magnetic field is important to the burst oscillation mechanism, either directly or by establishing composition or temperature gradients.  IGR~J17511--3057 has however also broadened the spectrum of properties exhibited by AMXP burst oscillations.  This AMXP is the first for which oscillation amplitudes drop below the detectability threshold in bright bursts, without the bursts exhibiting simultaneous PRE.  This phenomenon is also seen in non-pulsars, and points to a mechanism for suppressing burst oscillations - other than PRE - that is less affected by the presence of a dynamically important magnetic field.    \\textbf{Acknowledgment:} We thank M. Nowak for sharing the X-ray burst time of occurrence in their Chandra data, respectively. AP acknowledges support from the Netherlands Organization for Scientific Research (NWO) Veni Fellowship. ML acknowledges support from the Netherlands Organization for Scientific Research (NWO) Rubicon Fellowship."
    },
    "1005/1005.5250_arXiv.txt": {
        "abstract": "The characterization of the wind speed vertical distribution V(h) is fundamental for an astronomical site for many different reasons: (1) the wind speed shear contributes to trigger optical turbulence in the whole troposphere, (2) a few of the astroclimatic parameters such as the wavefront coherence time ($\\tau_0$) depends directly on V(h), (3) the equivalent velocity $V_0$, controlling the frequency at which the adaptive optics systems have to run to work properly, depends on the vertical distribution of the wind speed and optical turbulence. Also, a too strong wind speed near the ground can introduce vibrations in the telescope structures. The wind speed at a precise pressure (200 hPa) has frequently been used to retrieve indications concerning the $\\tau_0$ and the frequency limits imposed to all instrumentation based on adaptive optics systems, but more recently it has been proved that $V_{200}$ (wind speed at 200 hPa) alone is not sufficient to provide exhaustive elements concerning this topic and that the vertical distribution of the wind speed is necessary. In this paper a complete characterization of the vertical distribution of wind speed strength is done above Mt.Graham (Arizona, US), site of the Large Binocular Telescope. We provide a climatological study extended over 10 years using the operational analyses from the European Centre for Medium-Range Weather Forecasts (ECMWF), we prove that this is representative of the wind speed vertical distribution at Mt. Graham with exception of the boundary layer and we prove that a mesoscale model can provide reliable nightly estimates of V(h) above this astronomical site from the ground up to the top of the atmosphere ($\\sim$ 20 km). ",
        "introduction": "The characterization of the wind speed above an astronomical site is extremely important for several different reasons. Firstly, because the wind speed is strictly correlated to the optical turbulence: strong wind speed and sharp wind speed gradients are indicators of a turbulent atmosphere, which in combination with a stable stratification of the atmosphere (a positive gradient of the potential temperature) creates optical turbulence that limits the resolution of the telescopes.  Secondly, because the stronger the wind speed, the higher is the speed at which the turbulence layers cross the pupil of the telescope and the higher is the frequency at which adaptive optics systems are forced to work to correct the turbulence perturbations above the wavefront. Thirdly, because a too strong wind speed near the ground can introduce vibrations of the telescope structures.  At mid-latitudes, the wind speed varies with height showing a maximum at the jet stream level, usually 10-12 km (or $\\sim$200 hPa) above sea level. The strength of the wind speed at high altitudes varies according to season, with the strongest values during (the local) winter and early spring \\citep*{Ma01, Ca05, GL05, Ma06, Eg07, Bo09, Ma10}.  It has been observed that seasonal variations at the jet stream level are more pronounced at higher latitudes than at low latitudes in proximity of the equator \\citep{Ca05} because of the general circulation of the wind speed at synoptic scale.  At extreme latitudes, i.e. in proximity of the poles, the wind speed is characterized by a completely different feature in the free atmosphere \\citep{Ge06, Sa06}. % It is rather weak up to 10 km and then increases monotonically above this height with a rate that increases with the distance of the site from the centre of the polar vortex \\citep{Ha08}. % This assumption has been confirmed by \\cite{La09}. %  Apart from the fact that a windy atmosphere more easily triggers optical turbulence, the astroclimatic parameter that directly depends not only on the optical turbulence but also on the wind speed is the wavefront coherence time: \\begin{equation} \\label{tau} \\tau_0 = 0.057 \\lambda^{6/5} \\Big(\\int\\limits_{0}^{\\infty }{} V(h)^{5/3} C_N^2(h) dh \\Big)^{-3/5} \\end{equation} also equivalent to: \\begin{equation} \\tau_{0}=0.31\\frac{r_{0}}{V_{0}} \\end{equation} where $r_{0}$ is the Fried's parameter and V$_{0}$ is the equivalent velocity: \\begin{equation} V_{0}=\\left( \\frac{\\int\\limits_{0}^{\\infty }{}V(h)^{5/3} C_{N}^{2}\\left( h\\right) dh}{\\int\\limits_{0}^{\\infty }C_{N}^{2}dh}\\right) ^{3/5} \\label{vel} \\end{equation}  $\\tau_{0}$ depends, therefore, on the vertical distribution of the wind speed V(h) and the optical turbulence $\\CN2$(h). To simplify the calculation of the wavefront coherence time Sarazin \\& Tokovinin (2002) proposed a method to calculate $\\tau_{0}$ using only the wind speed estimated at 200 hPa instead of the whole profile using an empirical relationship between $V_{0}$ and $V_{200}$: firstly they used $V_0=0.4*V_{200}$ and, in a second time, the expression has been modified in $V_0=max(V_{ground},0.4*V_{200})$.  However, a few years ago \\citep{Ma06}, in a study done above San Pedro M\\'artir, it has been observed that the value of the constant 0.4 was not universal.  The authors found the value 0.56 above San Pedro M\\'artir. However, they also proved that the relative error introduced in $\\tau_{0}$ with the method of \\cite{Sa02}, that does not consider the vertical distribution of the wind speed but just the wind speed in two precise regions of the atmosphere (near the ground and at 200 hPa), could be as great as 20$\\%$-50$\\%$, even using the appropriated constant. In other words, they proved that the proportionality between V$_{0}$ and V$_{200}$ is poorly reliable if one wishes precise estimates of  $\\tau_{0}$ (even if we select an appropriated constant)\\footnote{More recently, other authors \\citep{GL09} calculated $V_0=max(V_{ground},0.4*V_{200})$ above the Teide Observatory and they found a third different value of the constant. However, we think they misunderstood the \\cite{Ma06} thesis because they stated that the \\cite{Sa02} method needs to be calibrated to be used.}. \\cite{Ma06} concluded therefore that, for the calculation of the $\\tau_{0}$, the vertical distribution of the wind speed on the whole troposphere is fundamental and necessary and the method suggested by \\cite{Sa02}  can provide only some qualitative (but not accurate) estimates because it presents some intrinsic weak points. It appears therefore absolutely important in the field of the site characterization of an astronomical site, at least in cases such as the calculation of $\\tau_{0}$, to characterize the wind speed vertical distribution. How to do that?  The estimate of the wind speed vertical distribution up to 20 km at an astronomical site (usually placed on the top of mountains) is not trivial.  The Generalized Scidar, an optical instrument based on a remote sensing principle, can measure the wind speed \\citep*{Av01} at all the heights in which turbulent layers are present i.e. it can reconstruct  a sort of vertical profile. Some estimates have been done in the past above different sites \\citep{Av06, Eg07, Ma10}. However, this method requires the employment of an instrument that has to be placed at the focus of a telescope with a pupil size of at least 1.5 m and it can be performed for short periods (typically some tens of nights) related to dedicated experiments. Such an instrument is not suitable for routinely monitoring the wind speed or climatological studies. Alternatively wind speed profiles are routinely calculated by the General Circulation Models (GCM) (mainly from the ECMWF and NOAA/NCEP) and data-set can be retrieved in any site in the world with a horizontal resolution of 0.25 degrees (operational analyses) and 2.5 degrees (re-analyses)\\footnote{To obtain re-analyses the GCM models are re-run as to offer an equal horizontal and vertical resolution and the same model configuration for many years in the past. In this way, the GCM outputs referring to the most recent years and the GCM outputs referring to older periods are made uniform. In other words, the performances of the re-analyses are degraded with respect to operational analyses to preserve the temporal uniformity. On the other hand the data-set is uniform all along decades and this is the reason why re-analyses are used mainly for climatological studies.}. Both have been used for astronomical applications. More precisely, the ECMWF analyses have been used at mid-latitude sites \\citep{Ma01, Eg07, Ma10} as well as at extreme latitudes \\citep{Ge06, Ha08}. NOAA/NCEP re-analyses (vertical profiles) have been used at mid-latitude sites \\citep{Ca05, Av06, Bo09} and ERA-reanalyses of the ECMWF have been used at extreme latitudes \\citep{Sa06}. Model outputs showed good correlations with measurements in all cases in which this has been calculated. The unique problem with analyses and re-analyses is that these estimates are representative of the wind speed above the astronomical site but not in the the surface and boundary layer where the local orographic effects have a major effect on the wind speed. Above roughly 1-2 km from the ground, the wind speed is almost horizontally homogeneous and the vertical distribution is basically the same with respect to a horizontal extension of some tens of kilometers. Below this height the analyses from the GCMs are less representative of the wind speed because their horizontal resolution is too low to give an accurate description of the interaction of the atmospheric flow with the topography near the surface.  However, in the surface (lowest few tens of meters) and boundary layers (typically the first kilometer above the surface) it has been proved \\citep{Ma03} that mesoscale models, with a horizontal resolution of 1 km, can provide much better estimates than what analyses from the GCMs do above mid-latitude astronomical sites. In that paper the author proved that mesoscale models can provide estimates better correlated to measurements than analyses from GCM. It has also been proved that mesoscale models are able, contrary to the analyses from the GCMs, to discriminate the wind speed near the ground between two astronomical sites (Paranal and Maidanak) characterized by a median wind speeds that differ for 4-5 ms$^{-1}$. This study lets us think that mesoscale models could be a useful tool to reconstruct the wind speed vertical profile all along the 20 km from the ground.  They are supposed to be comparable in performances to the ECMWF models above 1-2 km from the ground and they are supposed to provide better performances of the wind speed in the first 1-2 km from the ground.  In this paper we provide a climatological characterization of the wind speed above 1 km from the ground on the time scale of ten years using analyses from the ECMWF general circulation models and we investigate the possibility to use a mesoscale model to systematically reconstruct a complete wind speed vertical profile extended on the whole 20 km above Mt.Graham (Arizona, US) the site of the Large Binocular Telescope (LBT).  In Section 2 we use the operational ECMWF analyses over 10 years, 1998-2007, to present a climatological estimate of the median monthly wind speed. The operational data was chosen because of their higher resolution (0.25 degrees) with respect to the re-analyses. We will first prove the reliability of the ECMWF analyses in these regions of the world comparing analyses with radiosoundings launched in from Tucson International Airport ( $\\sim$120 km from Mt. Graham).  In Section 3 we investigate the reliability of the wind speed vertical profiles retrieved from simulations with a mesoscale model (Meso-NH) in the high as well as in the low part of the atmosphere. The mesoscale model is run in a grid-nesting configuration covering a total surface of 800 km x 800 km and three imbricated models with 10, 2.5 and 0.5 km horizontal resolution (Table \\ref{tab1}). The innermost model covering a surface of 60 km x 60 km.  To estimate the model reliability we use measurements of the wind speed done by a Generalized Scidar run at the focus of the Vatican Advanced Technology Telescope (VATT) as well as by an anemometer located on the roof of the same telescope.  This section aims to evaluate the possibility to use a mesoscale model to characterize the vertical distribution of the wind speed extending from the ground up to 20 km above astronomical sites. This should certainly represent an extremely valuable tool to provide an exhaustive monitoring of the wind speed above an astronomical site.  \\begin{table} \\caption{Meso-NH model configuration. In the second column the  horizontal resolution $\\Delta$X, in the third column the number of grid points and in the fourth column the horizontal surface covered by the model domain.} \\begin{tabular}{cccc} \\hline & $\\Delta$X (km) & Grid Points & Surface (km)\\\\ \\hline model 1 & 10  & 80 x 80 & 800 x 800\\\\ model 2 & 2.5 & 64 x 64 & 160 x 160 \\\\ model 3 & 0.5 & 120 x 120 & 60 x 60\\\\ \\hline \\end{tabular} \\label{tab1} \\end{table}  Finally, in Section 4, we present the conclusion of this study. ",
        "conclusions": "In this paper we aimed to give a complete characterization of the wind speed vertical distribution at Mt. Graham (Arizona, US) for astronomical applications. The simplest way of retrieving a complete characterization of the wind speed profile (with exception of the surface and boundary layer) is using the data from General Circulation Models (GCMs). In this study we use the operational analyses from the ECMWF extracted from the grid point closest to the peak of Mt. Graham, with a 0.25 degrees resolution, to study the vertical wind speed distribution over 10 years (1998-2007).  We have verified that the wind speed profile retrieved from the operational analyses of the ECMWF model is consistent with what is obtained by the radiosoundings from the nearby Tucson International airport. We also proved that the wind speed in this region is homogeneous with respect to horizontal spatial scales of the order of some tens of kilometers above $\\sim$4 km from the sea level.  Having proved that the ECMWF analyses are reliable, we presented the monthly median wind speed extended on a 10 years time scale. The wind speed profiles are rather typical of a mid-latitude site with a pronounced wind speed maximum at the jet stream level (10-12 km above sea level) during most of the year. On the contrary, during the summer the maximum wind speed is located well into the stratosphere. The strongest variability (from different years) in the monthly median wind speed is found during spring and autumn.   For the same period we also provided the monthly mean wind speed values at 200 hPa, corresponding to the maximum wind speed values at the jet-stream height. The month with the strongest wind speed at 200 hPa is February (37.21 ms$^{-1}$), while the weakest wind speed is observed in July (11.17 ms$^{-1}$). Results indicate that the wind speed at the jet-stream level is very similar to what has been observed above the Observatory in San Pedro M\\'artir (Baja California) and it is consistent with the values observed above the best astronomical sites in the world.  Besides, we proved the reliability of a mesoscale model (Meso-Nh) in reconstructing the wind speed on the whole 20 km just above the summit of Mt. Graham included the boundary layer and the surface layer. We proved that the wind speed reconstructed by the model is very well correlated to the ECMWF analyses and it is also able to provide realistic profiles in the first kilometer from the ground.  The wind speed estimates from the mesoscale model have been compared to measurements from a Generalized Scidar and an anemometer located at 20 m from the ground on a sample of 39 nights.  Above 50 m the wind speed profiles reconstructed by the model match in a very satisfactory way ($\\Delta$V $\\le$ 1 ms$^{-1}$) with respect to the measured wind speed profiles.  Closer to the surface, just in proximity of the top of the trees, the wind speed estimated by the model is included in the range of values given by the anemometer and the Generalized Scidar and for this reason can be considered satisfactory. However, the dispersion between the anemometer and the GS measurements seems a little too large. This difference is highly probably due to the fact the anemometer measures a wind speed that is not completely in free air (the anemometer is placed beside the dome of the VATT) and below the top of the trees. On the other side, the GS is probably affected by the wind speed just above the top of the trees because of its finite vertical resolution. The qualitative trend of the wind speed observed all along the sample of 39 nights is however very well reconstructed by the model and in agreement with measurements. The model appears to reconstruct very well the wind speed behavior above and below 20 m.  This paper therefore validates the Meso-Nh model as a tool to predict the wind speed vertical profile V(h) from the ground up to 20 km above Mt. Graham for each night and, at present time, it appears as the unique method to systematically estimate the whole wind speed vertical profile above an astronomical Observatory. We remind that it has been proved (Masciadri et al. 1999a, Masciadri \\& Jabouille, 2001, Masciadri et al. 2004, Lascaux et al. 2010) that Meso-Nh can provide reliable $\\CN2$ profiles above an astronomical site and it appears therefore as an extremely useful tool for $\\tau_{0}$ estimates. Besides there are evidences that it would be very useful to supply the Mt. Graham Observatory with anemometers located at different heights below and above the top of the trees ($\\sim$ 20 m) because this should permit to provide a better constraints of the model for dedicated and more detailed applications."
    },
    "1005/1005.1473_arXiv.txt": {
        "abstract": "{} {We exploit the deepest existing far-infrared (FIR) data obtained so far by Herschel at 100 and 160 $\\mu$m in the GOODS-N, as part of the PACS Evolutionary Probe (PEP) survey, to derive for the first time the evolution of the rest-frame 60-$\\mu$m, 90-$\\mu$m, and total IR luminosity functions (LFs) of galaxies and AGNs from $z=0$ to unprecedented high redshifts ($z$$\\sim$2--3).} {The PEP LFs were computed using the 1/V$_{max}$ method. The FIR sources were classified by means of a detailed broad-band SED-fitting analysis and spectral characterisation. Based on the best-fit model results, k-correction and total IR (8--1000 $\\mu$m) luminosity were obtained for each source. LFs (monochromatic and total) were then derived for various IR populations separately in different redshift bins and compared to backward evolution model predictions.} {We detect strong evolution in the LF to at least $z$$\\sim$2. Objects with SEDs similar to local spiral galaxies are the major contributors to the star formation density (SFD) at $z$$\\lsimeq$0.3, then, as redshift increases, moderate SF galaxies -- most likely containing a low-luminosity AGN -- start dominating up to $z$$\\simeq$1.5. At $z$$\\gsimeq$1.5 the SFD is dominated by the contributions of starburst galaxies. In agreement with previous findings, the comoving IR LD derived from our data evolves approximately as $(1+z)^{3.8\\pm\ufffd0.3}$ up to $z$$\\sim$1, there being some evidence of flattening up to $z$$\\sim$2.} {} ",
        "introduction": "Extragalactic surveys in the infrared (IR) represent a key tool for understanding galaxy formation and evolution, the luminosity function (LF) computed at different redshifts being the most direct method for exploring the evolution of a galaxy population. Deep surveys with {\\em IRAS} and {\\em ISO} allowed the first studies of the IR-galaxy LF at $z$$\\lsimeq$0.3 (Saunders et al. 1990) and at $z$$\\lsimeq$1, respectively (Pozzi et al. 2004). With {\\em Spitzer} 24-$\\mu$m data, it was possible to study the evolution of the mid-infrared (MIR) LF up to $z$$\\sim$2.5 (e.g. Le Floc'h et al. 2005; Caputi et al. 2007; Rodighiero et al. 2010), while, even with the deepest {\\em Spitzer} 70-$\\mu$m data, only $z$$\\sim$1 could be reached in the far-infrared (FIR) (Magnelli et al. 2009). Ground-based observations in the mm/sub-mm range, probing the evolution of the most distant ($z$$\\gsimeq$2) and luminous dusty galaxies, have so far been limited to the identification of sources at the very bright end of the luminosity function (e.g., Chapman et al. 2005). All of these works detected strong evolution in both luminosity and density, indicating that IR galaxies were more numerous and more luminous in the past. The derived rates of evolution of the IR sources are significantly higher than observed at any other wavelength. Since the rest-frame IR spectral energy distributions (SED) of star-forming galaxies and AGNs peak at 60--200 $\\mu$m, to measure their bolometric luminosity and evolution with $z$ we need to observe in the FIR/sub-mm regime. However, the detection of large numbers of high-$z$ sources at the peak of their IR SED could not yet been achievable, due to source confusion and/or low detector sensitivity.\\\\ \\indent In the present paper, we discuss observations performed with the PACS instrument (Poglitsch et al. 2010) onboard the {\\em Herschel Space Observatory} (Pilbratt et al. 2010). {\\em Herschel} is the first telescope allowing us to detect the FIR population to high redshifts ($z$$\\sim$3-4) and to derive its rate of evolution by a detailed LF analysis. The PACS Evolutionary Probe (PEP, PI D. Lutz) Surveys are designed specifically to this purpose. With the observations dedicated to the GOODS-N  by the PEP Science Demonstration Program (SDP), we first address the main goal of determining the evolution with redshift of the galaxy and AGN LF in the FIR domain. The availability of a deep multiwavelength catalogue in the GOODS-N is crucial for analysing the SEDs, obtaining k-corrections and total IR luminosities, and studying the LF and its evolution separately for the different IR populations. The measure of the total IR luminosity allows us to derive for the first time the total IR LF and its evolution directly from FIR data (see e.g. Nordon et al. (2010), this issue, for descriptions of the failure, at least at $z$$>$1.5, of previous total IR luminosity extrapolations from the MIR).\\\\ \\indent Throughout this paper, we assume a $\\Lambda$CDM cosmology with H$_0$\\,=\\,71~km~s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_m$\\,=\\,0.27, and $\\Omega_{\\lambda}\\,=\\,0.73$. ",
        "conclusions": "We have used of the deepest {\\em Herschel} 100- and 160-$\\mu$m data from the PEP SDP survey in the GOODS-N to characterise the evolution of the galaxy and AGN FIR LF and LD across the redshift range 0$\\lsimeq$$z$$\\lsimeq$3. To summarise, in the present work we have: \\begin{enumerate} \\item completely characterised the multiwavelength SEDs of the PEP sources by performing a detailed SED-fitting analysis, % tightly constraining (and in some cases opportunely modifying) the templates at long wavelengths. \\item computed the first rest-frame LFs at 60 and 90 $\\mu$m up to $z$$\\sim$4, finding that a significant amount of evolution is required to reproduce the LF at least up to $z$$\\sim$2, the different SED-classes exhibiting different evolutionary behaviours. \\item integrated the SEDs at $\\lambda_{rest}$$=$8--1000 $\\mu$m and computed the total IR LF up to $z$$\\sim$3, % finding luminosity evolution $\\propto$$(1+z)^{4.1^{+0.3}_{-0.2}}$ up to $z$$=$$1.5\\pm0.3$, then an almost constant evolution rate to $z$$\\sim$2.5--3. \\item derived the evolution of the comoving total IR LD (and SFD), which was found to increase as $(1+z)^{3.8\\pm0.3}$ up to $z$$\\sim$1. Spiral galaxies are the main contributors to the SFD at $z$$\\lsimeq$0.3, moderate SF galaxies (which most likely harbour a low luminosity AGN) becoming increasingly important up to $z$$\\simeq$1.5. At this redshift, the SFD is representative of SB galaxies. \\end{enumerate}"
    },
    "1005/1005.3195_arXiv.txt": {
        "abstract": "We study the high-energy emission of the Galactic black hole candidate GX 339--4 using \\textit{INTEGRAL}/SPI and simultaneous \\textit{RXTE}/PCA data. By the end of January 2007, when it reached its peak luminosity in hard X-rays, the source was in a bright hard state. The SPI data from this period show a good signal to noise ratio, allowing a detailed study of the spectral energy distribution up to several hundred keV. As a main result, we report on the detection of a variable hard spectral feature ($\\geq$$150$\\,keV) which represents a significant excess with respect to the cutoff power law shape of the spectrum. The SPI data suggest that the intensity of this feature is positively correlated with the $25$\\thinspace --\\thinspace $50$\\,keV luminosity of the source and the associated variability time scale is shorter than $7$\\,hours. The simultaneous PCA data, however, show no significant change in the spectral shape, indicating that the source is not undergoing a canonical state transition. We analyzed the broad band spectra in the lights of several physical models, assuming different heating mechanisms and properties of the Comptonizing plasma. For the first time, we performed quantitative model fitting with the new versatile Comptonization code \\textsc{belm}% , accounting self-consistently for the presence of a magnetic field. We show that a magnetized medium subject to pure non-thermal electron acceleration provides a framework for a physically consistent interpretation of the observed $4$\\thinspace --\\thinspace $500$\\,keV emission. Moreover, we find that the spectral variability might be triggered by the variations of only one physical parameter, namely the magnetic field strength. Therefore, it appears that the magnetic field is likely to be a key parameter in the production of the Comptonized hard X-ray emission. ",
        "introduction": "GX 339--4 was discovered in the early 70's by the MIT X-ray detector aboard the OSO-7 mission \\citep{markert1973}. The source is classified as a low mass X-ray binary (LMXB; \\citealt{shahbaz2001}) from upper limits on the optical luminosity of the companion star. It is believed to harbor a black hole, for which \\citet{hynes2003} derived a mass function of $5.8 \\pm 0.5\\,\\rm M_{\\sun}$. The inclination of the system is yet uncertain. A study of the binary parameters \\citep{zdziarski2004} revealed a plausible lower limit of $i \\geq 45^{\\circ}$, while \\citet{cowley2002} suggested $i \\leq 60^{\\circ}$ because the system is not eclipsing. On the other hand, spectral fits of the Fe $K_{\\alpha}$ region with Chandra \\citep{miller2004a} and XMM Newton \\citep{miller2004b,reis2008} clearly favor lower values for the inner disc inclination (typically $i \\sim 20^{\\circ}$). This may indicate that the inner accretion disc is warped, making it appear at a lower inclination than the orbital plane. Similarly, there is no certainty concerning the distance to the source. Resolving the velocity structure along the line of sight, \\citet{hynes2004} obtained a conservative lower limit of $d \\geq 6$\\,kpc. \\citet{zdziarski2004} analyzed the binary parameters of the system and found the most plausible distance to be around $8$\\,kpc. In this work, we use $d = 8$\\,kpc, $i = 50^{\\circ}$ and, assuming a small companion mass, we infer $M=13\\,\\rm M_{\\sun}$ from the mass function.  Black hole binaries (BHBs) are powerful engines producing high-energy radiation up to the $\\gamma$-ray domain. From an observational perspective, they appear in four spectral states, namely, quiescent, low/hard, intermediate and high/soft (\\citealt{tanaka1995}, \\citealt{belloni2005}; see also \\citealt{mr06} for a slightly different classification).   {\\it Soft states:} In the soft state, the energy spectrum is dominated by a soft ($\\leq 10$\\,keV) component which is attributed to thermal emission from an optically thick geometrically thin accretion disc \\citep{ss73} extending down to the last stable orbit. Above $10$\\,keV, the emission is characterized by a complex hard X-ray continuum with no clear presence of a high-energy cutoff (see e.g. \\citet{gierlinski1999,motta2009,caballero-garcia2009}). This component can be attributed to inverse Compton scattering of soft photons (UV, soft X) in a hybrid thermal/non-thermal electron plasma (the so-called 'corona') (see e.g. the review by \\citealt*{dgk07}, hereafter DGK07).  The non-thermal particles in the corona may be generated due to the magnetic field. Indeed, the Parker instability is able to transport a significant fraction of the accretion power above and below the disc \\citep{galeev1979,uzdensky2008}, where the energy may then be dissipated through magnetic reconnection. Alternatively, Fermi acceleration at relativistic shocks is also expected to produce non-thermal particle distributions which can explain the observed steep power law emission in hard X-rays.  {\\it Hard states:} In the hard state, the energy spectrum is very different. The disc emission almost vanishes and the spectrum is dominated by a hard power law component (photon index $\\Gamma =$ 1.4 -- 2.0) with a nearly exponential cutoff ($E_{\\rm cut}$ = 50 -- 150\\,keV) (see e.g. \\citealt{zdziarski1998}). Such a spectrum is well described by thermal Comptonization in a hot, optically thin electron-proton plasma \\citep{sunyaev1979,zg04}. Soft $\\gamma$-ray observations of several BHBs additionally revealed a high-energy excess with respect to a thermal Comtonization model, suggesting that in some cases at least some level of non-thermal electron acceleration is also required. For instance, COMPTEL observed a non-thermal tail in the averaged hard state spectrum of Cygnus X-1 \\citep{mcconnell2002} while OSSE and SPI detected such a feature during bright hard states of GX 339-4 \\citep{johnson1993,wardzinski2002,joinet2007}.  {\\it Intermediate states:} Intermediate states are observed during transitions between the hard and the soft states and usually show the characteristic features of both \\citep{miyamoto1991,mendez1997,kong2002}. For GX 339--4, \\citet{belloni2005} defined two different varieties of intermediate state: the hard intermediate state (HIMS) and the soft intermediate state (SIMS). The transition from the HIMS to the SIMS of the 2004 outburst of GX 339--4 is reported in \\citet{delsanto2008}.  According to a popular scenario, the different spectral states can be explained through changes in the geometry of the accretion flow. The weakness of the thermal component in the hard state is generally interpreted as a consequence of a truncated accretion disc (DGK07), which, in its inner parts, is replaced by a hot, advection dominated accretion flow (ADAF; \\citealt{shapiro1976,narayan1994,yuan2007}). In these solutions, gravitational energy is converted into thermal energy of protons, which in turn heat the electrons through Coulomb collisions. This process naturally forms the quasi thermal electron distributions that are required to explain the typical hard state spectra. Moreover, recent models of hot accretion flows include a small non-thermal component that can account for the sometimes observed high-energy excess.  However, a number of recent results seem to question this paradigm. First, it is notoriously difficult to estimate the inner disc radius in the hard state and its recession is a highly debated topic \\citep[see e.g.][]{cabanac2009}. Although several reports support the truncated disc scenario \\citep{tomsick2009,done2009}, others suggest that even in hard states the accretion disc may extend down to a few Schwarzschild radii \\citep{miller2006,reis2008}, which would be inconsistent with the standard ADAF scenario. Second, hot accretion flows generally exist only if the optical depth is small ($\\tau_T \\ll1$), which implies that the electron-proton coupling is weak and therefore allows high proton temperatures. However, it is difficult to align such a small optical depth with both the hard X-ray spectral slope and the thermal cutoff energy. When the spectral slope is reproduced, classic ADAF models lead to higher electron temperatures than those observed in BHBs, as it was shown for Cygnus X-1 or GX 339-4 \\citep[hereafter MB09]{yuan2004,mb09}. Recent Monte-Carlo simulations, which account for global Compton cooling and general relativistic effects \\citep{xie2010}, yield self-consistent solutions in which the peak energy of the spectrum is reduced with respect to previous results. However, it is still not straightforward to accommodate the inferred high-energy cutoff to the observations, even if gravitational redshift is accounted for.  As an alternative to the ADAF-like models, the Comptonizing medium in the hard state could as well be powered by the same non-thermal mechanisms that are believed to accelerate the electrons in the soft state, i.e. diffusive shock acceleration or magnetic reconnection. Such a model naturally accounts for the presence of a non-thermal component in the high-energy spectrum. In addition, MB09 showed that the steady state electron distribution can appear quasi thermal even if acceleration mechanisms are purely non-thermal (see also \\citet{poutanen2009}). These authors studied the thermalizing effects of the magnetic field, since it was pointed out by \\citet*{ghisellini1998} that the very fast emission and absorption of synchrotron photons (the so-called 'synchrotron boiler effect') is able to thermalize the electron distribution in a few light-crossing times. Using \\textsc{belm}, a new code which includes this effect \\citep*{bmm08}, MB09 qualitatively explained the variety of spectral states observed in the prototypical BHB Cygnus X-1. The model is consistent with a disc recession in the hard state \\citep{poutanen2009}, but since the weakness of the soft component could as well result from a lower disc temperature \\citep{beloborodov1999,malzac2001}, it does not necessarily require a change in the geometry of the accretion flow (MB09).  In this paper, we use observations of GX 339--4 with SPI/\\textit{INTEGRAL} and PCA/\\textit{RXTE} to study the spectral energy distribution of the bright hard state from the 2007 outburst. We analyze the spectral variability in the framework of the models cited herebefore. The data and the associated reduction procedures are described in Section 2. Section 3 is dedicated to a phenomenological analysis of the high-energy behavior while in Section 4 we present an extensive physical analysis of the broad band spectra. The results and their implications are discussed in Section 5 and summarized in the conclusive section of the paper. ",
        "conclusions": "We presented an analysis of the high-energy emission of GX 339--4 in a luminous hard state. With respect to the standard cutoff shape of the hard state spectrum, the 25 -- 500\\,keV \\textit{INTEGRAL}/SPI data revealed the appearance of a variable high-energy excess. The intensity of this hard excess seems to be positively correlated with the total X-ray luminosity and the associated time scale is shorter than $7$\\,hours. We explored the possible physical origins of this variability through an extensive analysis of two averaged spectra, one showing the typical cutoff shape and one showing this prominent high-energy excess.  We used simultaneous \\textit{RXTE}/PCA data to extend the spectral coverage down to $4$\\,keV and fitted the broad band spectra with a variety of physical Comptonization models. Models involving only thermal heating can be ruled out since they are not able to reproduce the high-energy tail. This feature thus confirms that in luminous hard states, the Comptonizing plasma of GX 339--4 contains a fraction of non-thermal particles. Models involving only non-thermal electron acceleration, on the other hand, showed that the thermal part of the spectrum can be produced by an initially non-thermal (power law) distribution, which rapidly thermalizes under the effects of synchrotron self-absorption and/or \\textit{e-e} Coulomb collisions.  The relatively good signal to noise ratio of the high-energy channels ($>$$150$\\,keV) allowed to derive meaningful constraints to the model parameters. Depending on the nature of the plasma (magnetized or not), the transition between the two averaged spectra requires the variations of at least two parameters. We found that a magnetized medium subject to non-thermal electron acceleration provides the framework for a simple and physically consistent interpretation of the data. Indeed, we showed that the spectral variability can be triggered by the \\textit{only} variations of the magnetic field, implying a subsequent variation of the maximum energy of the accelerated particles.  The quantitative constraints derived from this model yield a very conservative upper limit on the average magnetic field strength in the Comptonizing plasma and suggest that in the bright hard state, the magnetic energy density could reach equipartition with the radiative energy density.  In conclusion, the presented results suggest that magnetic processes are likely to play a crucial role in the production of the high-energy emission of GX 339--4. In luminous hard states, the Comptonized emission could originate from a magnetized corona essentially powered though non-thermal particle acceleration, similarily to what is believed to happen in soft states. The hard X-ray emission in both spectral states may therefore be the consequence of a common physical phenomenon."
    },
    "1005/1005.4126_arXiv.txt": {
        "abstract": "The effect of stellar scintillation on the accuracy of photometric measurements is analyzed. We obtain a convenient form of estimaton of this effect in the long exposure regime, when the turbulence shift produced by the wind is much larger than the aperture of the telescope. A simple method is proposed to determine index $S_3$ introduced by \\citet{ken2006}, directly from the measurements with the Multi Aperture Scintillation Sensor (MASS) without information on vertical profile of the wind. The statistics $S_3$ resulting from our campaign of 2005 -- 2007 at Maidanak observatory is presented. It is shown that these data can be used to estimate high-altitude winds at pressure level 70 -- 100 mbar. Comparison with the wind speed retrieved from the NCEP/NCAR global models shows a good agreement. Some prospects for retrieval of the wind speed profile from the MASS measurements are outlined. ",
        "introduction": "The accuracy of measurements of brightness of astronomical objects is affected by many factors. It is clear that they are influenced in different ways depending on the applied photometric techniques. One of these effects is the stellar scintillation caused by {\\it optical turbulence} (OT) in the Earth's atmosphere. Unlike other variations of the atmospheric transparency, the light fluctuations due to the scintillation do not correlate for the images separated by more than the isoplanatic angle ($2''-5''$ in the optical range). Therefore, errors introduced by this effect appear always and increase when the brightness of two stars is compared.  This problem was noted and studied a long time ago \\citep{Young69, DravI, DravII, DravIII}. Recently, interest in this problem manifested itself again, but in the overall context of a study and characterization of the OT in various astronomical observatories and prospective sites. In the paper \\citet{ken2006} a special index $S_3$ have been introduced to relate the vertical distribution of the turbulence and wind speed to the variations of the light flux as measured in a large telescope.  Monitoring of this parameter and its analysis is very useful for comparing different sites and for predicting or planning of photometric observations. However, the approach used in cited work implies prior restoration of the OT profiles from the MASS -- Multi Aperture Scintillation Sensor \\citep{MASS, mnras2003} data and the use of additional information on wind speed. This is a very intensive procedure, besides burdened by additional errors  In this paper we propose a method to determine the index $ S_3 $ directly from the MASS measurements without using additional information. The presented method is preceded by a theoretical consideration of the scintillation in the temporal domain. ",
        "conclusions": "A method to obtain statistical prediction of photometric accuracy based only on the data from the MASS instrument is proposed. Comparison of the results obtained by this method and by integration of the OT and wind profiles \\citep{ken2006} can be used to verify wind models above astronomical sites.  It is particularly interesting to perform such a comparison for the measurements with MASS instrument performed in Antarctica at Dome C, as described in \\citep{ken2006}. There is some concern that their data censoring procedure, leaving only 16\\% of useful data, may bias the $S_3$ estimates. An ability to estimate wind speed at pressure level of 70 -- 100~mbar directly from the MASS data will help to check the applicability of the wind models and will avoid biases caused by replacing instantaneous values by their statistical estimates.  Additionally, we note here that 1-s averaging time in the case of MASS apertures is too long. A regime of the long exposure with the relevant asymptotics starts at shorter exposures on the order of 0.1 --- 0.2~s, and the use of such short time would markedly improve the accuracy and reliability of the temporal index. The behavior of WFs $U'(h)$ (Fig.~\\ref{fig:up}) indicates that by using a set of precise measurements of temporal indices and the restored $C_n^2(h)$ profile, a reasonable solution of the system equations like (\\ref{eq:long1}) for the unknowns $1/w(h)$ may be achieved.  Realization of this possibility will permit 1) to obtain additional information about the current state of the turbulent atmosphere, and 2) to correct more precisely the MASS/DIMM data for the effects of finite exposure time.  The author thanks his colleagues who participated in the campaign to measure the optical turbulence above the mountain Maidanak in 2005 -- 2007, producing the data used here. The work was partially supported by the RFBR grant 06-02-16902a."
    },
    "1005/1005.2518.txt": {
        "abstract": "In loop quantum cosmology, the Hamiltonian reduces to a finite difference operator. We study the initial singularity and the large volume limit against the ambiguities in the discretisation and the operator ordering within a homogeneous, isotropic and spatially flat model with the cosmological constant. We find that the absence of the singularity strongly depends on the choice of the operator ordering and the requirement for the absence singles out a very small class of orderings. Moreover we find a general ordering rule required for the absence of the singularity. We also find that the large volume limit naturally recovers a smooth wave function in the discretisation where each step corresponds to a fixed volume increment but not in the one where each step corresponds to a fixed area increment. If loop quantum cosmology is to be a phenomenological realisation of full loop quantum gravity, these results are important to fix the theoretical ambiguities. % We study the flat Friedmann-Robertson-Walker universe with the cosmological constant % in the context of loop quantum cosmology. % If one considers the original scheme used in Ref.~\\cite{Bojowald4}, % no continuum limit problem arises in the model. % We show that the problem is resolved by using the improved scheme introduced in Ref.~\\cite{O18_Ashtekar}. % It is also shown that the initial singularity is avoided in the improved scheme only when an appropriate operator ordering of the Hamiltonian constraint is chosen. ",
        "introduction": "Loop quantum gravity (LQG) is a background-independent nonperturbative quantum gravity. %LQG is a quantisation of There, one quantises the Hamiltonian formulation of general relativity based on the SU(2) connection $A^{i}_{a}$ and the densitised triad $E^{a}_{i}$ on the three-dimensional space, where $a$, $b$, ... and $i$, $j$, ... both run over 1, 2 and 3. $a$, $b$, ... are tangent indices, while $i$, $j$, ... are indices associated with a basis $\\{\\tau_{i}\\}$ for the Lie algebra of SU(2) Lie group and raised and lowered with Kronecker's delta $\\delta_{ij}$. The three-dimensional inverse metric is given by $q^{ab}=|\\det E^{c}_{k}|^{-1}E^{a}_{i}E^{b}_{j}\\delta^{ij}$. To quantise the system, LQG invokes holonomies $h_{e}=e^{\\int_{e}A^{i}_{a}\\tau_{i}dx^{a}}$ along an edge $e$ and fluxes ${\\cal E}_{\\cal S}=\\int_{\\cal S} d^{2}\\sigma n_{a} E^{a}_{i}\\tau^{i}$ over a two-surface ${\\cal S}$, where $d^{2}\\sigma$ is the area element on ${\\cal S}$ and $n_{a}$ is a unit normal to ${\\cal S}$. There appears a nondimensional constant parameter $\\gamma$, which is called the Barbero-Immirzi parameter %~\\cite{Barbero} and cannot be determined within the theory. The kinematical Hilbert space is spanned by the spin network states. The Hamiltonian as well as the area and the volume are constructed from holonomies and fluxes and act as operators on the kinematical Hilbert space. The Hilbert space is spanned by equivalence classes of the spin networks under diffeomorphisms or the s-knot states. See~\\cite{Rovelli0, Thiemann0, Ashtekar_2005} for the basics and recent developments in LQG.  Loop quantum cosmology (LQC)~\\cite{lqc} is a quantised minisuperspace model motivated by LQG. In traditional quantum cosmology~\\cite{traditional_QC}, symmetry reduced models are quantised in the usual Schr\\\"{o}dinger representation and the quantised Hamiltonian constraint yields the Wheeler-DeWitt (WDW) differential equation, while in LQC one uses the so-called polymer particle representation~\\cite{polymer} which is unitary inequivalent to the Schr\\\"{o}dinger representation and obtains a second order {\\it difference equation} rather than the differential equation. Remarkably, Bojowald~\\cite{Bojowald4} demonstrated that there can be no big bang singularity in the following two aspects. First, %there appears no curvature singularity because the spectrum of the inverse scale factor operator is bounded from above. Second, the wave function of the universe can be uniquely extended beyond the point which was the initial singularity in classical theory. The above features have been first shown for a homogeneous, isotropic and spatially flat model and subsequently generalised to nonflat or anisotropic cases~\\cite{lqc}, although it is yet uncertain whether this can be generalised to inhomogeneous models %although it is yet uncertain %whether this generalises to inhomogeneous models (see e.g.~\\cite{bojowald_etal_2008,bojowald_etal_2009}).  In this paper, we visit the absence of the singularity in the presence of the cosmological constant against the quantisation ambiguities. In particular, we focus on the choice of the discretisation parameter $\\lambda$ which corresponds to the coordinate length of edges for the basic holonomies and the operator ordering in the Hamiltonian constraint. To study these issues, we fix the cosmological model to be homogeneous, isotropic and spatially flat with the cosmological constant. We first adopt the constant $\\lambda$ as in~\\cite{Bojowald4}, for which the universe gains a fixed area quantum at each step and find that the absence of the singularity strongly depends on the operator ordering in the Hamiltonian. However, for this discretisation, we encounter a serious problem in the physical interpretation of the obtained wave function in the large volume limit. The problem is that the behaviour of the wave function obtained from the difference equation in LQC does not agree with that obtained from the WDW equation in the large volume limit. Since the effects of discreteness are dominant in the Planck scale physics but should disappear when the universe becomes large, the discrete wave function in LQC in the large volume is expected to have a correspondence to the smooth wave function which is a solution to the WDW equation. We see that this problem is resolved if $\\lambda$ is chosen to vary~\\cite{Ashtekar} so that the universe gains a fixed volume quantum and find that the absence of the singularity strongly depends on the operator ordering in the Hamiltonian also in this discretisation. It is already studied in detail how the large volume limit depends on the choice of the discretisation in the presence of matter fields by Nelson and Sakellariadou~\\cite{Nelson_Sakellariadou_2007a,Nelson_Sakellariadou_2007b} and the present result with the cosmological constant but without matter fields is consistent with theirs.  This paper is organised as follows. In Section~\\ref{sec:lqc}, we introduce the discretised Hamiltonian for a homogeneous, isotropic and spatially flat model with the cosmological constant. In Section~\\ref{sec:equi-area}, we choose $\\lambda$ to be constant and derive the difference equation. Then, we demonstrate how the absence of the singularity depends on the choice of the operator ordering and numerically and analytically show the large volume limit is problematic. In Section~\\ref{sec:equi-volume}, we vary $\\lambda$ appropriately and derive the difference equation. Then, we show that the large volume limit problem is resolved in this discretisation and also find that the absence of the singularity strongly depends on the choice of the operator ordering. Section~\\ref{sec:discussion} is devoted to discussion for the influence of the matter fields. %for the singularity avoidance. %is devoted to discussion. In Section~\\ref{sec:conclusion} we conclude the paper. In this paper we use the units in which $c=\\hbar=1$. ",
        "conclusions": "\\label{sec:conclusion} We have investigated a homogeneous, isotropic and spatially flat universe with the cosmological constant in the context of LQC. In particular, we have studied theoretical ambiguities in the quantisation of the Hamiltonian, which arise in the operator ordering and the discretisation and are hard to fix within the framework of LQC. We focus on two important features of LQC, the absence of the initial singularity and the large volume limit. We have shown that the absence of the initial singularity strongly depends on the choice of the operator ordering. Therefore, the requirement for the absence of the singularity can potentially make the ordering ambiguity very small. Furthermore, we have found a general ordering rule required for the absence of the singularity. On the other hand, the choice of the discretisation is crucial when we consider the large volume limit. We have demonstrated two typical choices of the discretisation, the equi-area and equi-volume ones. Then, we have shown that in the former case there arises a serious problem in the physical interpretation of the wave function as a continuous function in the large volume limit and also that this problem is resolved in the latter case. It is clear that we cannot fix these ambiguities only from the present results. However, if we can deduce the physical implications of full LQG even qualitatively, we will be able to fix the ambiguities of LQC as the phenomenological realisation of full LQG in the minisuperspace model.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%% ACKNOWLEDGEMENT %%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1005/1005.2929.txt": {
        "abstract": "A small percentage of normal stars harbor giant planets that orbit within a few tenths of an astronomical unit. At such distances the potential exists for significant tidal and magnetic field interaction resulting in energy dissipation that may manifest as changes within the stellar corona. We examine the X-ray emission of stars hosting planets and find a positive correlation between X-ray luminosity and the projected mass of the most closely orbiting exoplanets. We investigate possible systematics and observational biases that could mimic or confuse this correlation but find no strong evidence for any, especially for planets more massive than $\\sim 0.1$ M$_J$. Luminosities and upper limits are consistent with the interpretation that there is a lower floor to stellar X-ray emission dependent on close-in planetary mass. Under the hypothesis that this is a consequence of planet-star magnetic field interaction, and energy dissipation, we estimate a possible field strength increase between planets of 1 and 10 M$_J$ of a factor $\\sim 8$. Intriguingly, this is consistent with recent geodynamo scaling law predictions. The high-energy photon emission of planet-star systems may therefore provide unique access to the detailed magnetic, and hence geodynamic, properties of exoplanets. ",
        "introduction": "Surveys for exoplanets have revealed a population of objects that orbit with very short periods (e.g. \\citet{marcy2005}). The potential therefore exists for significant tidal and magnetic interaction between these planets and their host stars \\citep{cuntz2000,ip2004,preusse2006,mcivor2006,cranmer2007}.  Tidal evolution of planetary orbits, while still subject to significant uncertainty (due to the poorly constrained tidal quality factor $Q_p$, e.g. Matsumura et al. 2008), is well established (e.g. Jackson et al. 2008 and references therein). Tidal dissipation within planetary atmospheres is also a potentially important energy source in models of gas giant characteristics (e.g. Ibgui et al. 2009). In many cases the tidal evolution is dominated by the torques on the bulge raised on a planet by its parent star. However, for the most massive and closest orbiting planets, then significant torques are also possible due to the bulges raised in the stellar atmosphere. These modify the orbital evolution of the planet and act (for prograde planetary orbits) to maintain a higher observed spin rate for the host stars at a given age. Such an effect has been claimed in studies of exoplanet host star rotational velocities (Pont 2009), enhanced UV luminosities (Shkolnik et al. 2008) and at least one system - $\\tau$-Bootis - shows evidence for stellar synchronicity with a planetary orbit \\citep{walker2008}. It is also well established that main-sequence stellar X-ray luminosity is correlated with stellar rotation rates \\citep{ayres1980,pallavicini1981,maggio1987} (see below), which are generally higher for the youngest objects.  Magnetic field interaction between stars and planets is likely complex, but is expected to exhibit distinctly different characteristics to tidally induced phenomena. In particular (and see below) we might expect stellar variations directly related to planetary orbital phase and correlation with planetary mass and orbital parameters. There is growing evidence from the study of stellar photometry \\citep{walker2008}, chromospheric emission \\citep{shkolnik2008}, and X-ray emission \\citep{kashyap2008, poppenhaeger2010} that interaction of this nature is indeed occurring between stars and planets - at least in certain cases. Furthermore, although the physical origin is not clear, \\citet{hartman2010} finds a robust correlation between stellar chromospheric activity and the surface gravity of transiting hot Jupiters.  X-ray, emission from the stellar corona is a particularly interesting probe of star-planet interaction. Coronal X-ray photons originate from thermal bremsstrahlung, and inner atomic transitions of the hot, $10^6-10^7$ K, plasma. The general heating mechanisms in stellar coronae are still largely unknown, but appear to be intimately related to magnetic field structures and complex transferral of energy from the stellar surface \\citep{aschwanden2004,jess2009}. The placement of a planet, particularly a gas giant of Jupiter mass or greater, in proximity to the stellar coronal magnetic field may  effect the transport and release of energy in the plasma and therefore the high-energy photon output. Indeed, if stellar coronal magnetic field alignment favors strong coupling with planetary magnetospheres then significant, time-dependent, energy release may occur. Estimates suggest excess X-ray emission in the range of of $\\sim 10^{27}$ erg s$^{-1}$ to as much as an order of magnitude variation in the quiescent stellar X-ray luminosity is possible \\citep{ip2004,lanza2008,lanza2009b,cohen2009}).  Stellar X-ray emission is of course determined by a variety of systematic and stochastic components, related to fundamental stellar activity (see above). The addition of an emission component that may vary on timescales commensurate with a planetary orbit poses a challenge for observations. At the distances for which both magnetic and tidal interactions are expected to become significant ($\\sim 0.1-0.2$ AU) then orbital periods will range from $\\sim 1$ to $\\sim 20$ days (assuming stellar masses close to 1 M$_{\\odot}$). Many X-ray pointed observations of nearby stars span only a few hours and may not adequately sample the possible planet-induced emission variation. Current techniques of planet detection \\citep{otoole2009} are also significantly dependent on overall stellar activity, introducing further observational biases in sample selection.  In order to further investigate the  possibility of coronal emission signatures correlated with close planetary companions we use a sample of exoplanet host stars within 60 pc with both X-ray detections and upper limits from the R{\\\"o}ntgensatellit (ROSAT) mission all-sky survey. We describe this data in \\S2 and discuss its unique, partially time-averaged, characteristics which are well suited to this problem.  We examine the detailed relationship of X-ray luminosities and upper limits to planetary properties (\\S 2.1, 3), and examine potential sources of bias (\\S4). In \\S 5 we consider tidal effects, and in \\S6 we consider planet-star magnetic field interaction and the possible interpretation of our results. In \\S 7 we discuss and summarize our results and conclude in \\S 8. ",
        "conclusions": "The data presented here indicate that X-ray emission from planet-hosting stars may offer a probe of planet-star interactions, and conceivably a probe of the inner workings of close-orbit giant worlds through their convective-dynamo generated magnetic fields. The present sample is however small and more sensitive, targeted, X-ray data on planet hosting stars is needed. This data must also fully sample emission during the entire planetary orbital period in order to avoid biasing any measurement. Current observational programs are beginning to fill this gap (e.g. \\citet{poppenhaeger2010, sanz-forcada2010}), but sample selection must also be made in a uniform and unbiased fashion in order for these data to be of greatest utility. Given the inherent biases in RV and transit detection of planets this will be particularly critical.  As present works show \\citep{shkolnik2008,walker2008}, it is clear that planet-star interaction in specific cases can directly affect observables at a variety of wavelengths, and that the detailed study of individual objects will also be extremely valuable. In particular, obtaining sensitive time-series X-ray data of close-orbit, {\\em transiting} planet systems over multiple orbital periods and epochs may help disentangle inclination effects, magnetic field dissipation or plasma density enhancement mechanisms, and intrinsic stellar variations.   %% In a manner similar to \\objectname authors can provide links to dataset %% hosted at participating data centers via the \\dataset{} command.  The %% second curly bracket argument is printed in the text while the first %% parentheses argument serves as the valid data set identifier.  Large %% lists of data set are best provided in a table (see Table 3 for an example). %% Valid data set identifiers should be obtained from the data center that %% is currently hosting the data. %% %% Note that AASTeX interprets everything between the curly braces in the %% macro as regular text, so any special characters, e.g. \"#\" or \"_,\" must be %% preceded by a backslash. Otherwise, you will get a LaTeX error when you %% compile your manuscript.  Special characters do not %% need to be escaped in the optional, square-bracket argument.   %% In this section, we use  the \\subsection command to set off %% a subsection.  \\footnote is used to insert a footnote to the text.  %% Observe the use of the LaTeX \\label %% command after the \\subsection to give a symbolic KEY to the %% subsection for cross-referencing in a \\ref command. %% You can use LaTeX's \\ref and \\label commands to keep track of %% cross-references to sections, equations, tables, and figures. %% That way, if you change the order of any elements, LaTeX will %% automatically renumber them.  %% This section also includes several of the displayed math environments %% mentioned in the Author Guide.    %% The equation environment wil produce a numbered display equation.   %% The \\notetoeditor{TEXT} command allows the author to communicate %% information to the copy editor.  This information will appear as a %% footnote on the printed copy for the manuscript style file.  Nothing will %% appear on the printed copy if the preprint or %% preprint2 style files are used.  %% The eqnarray environment produces multi-line display math. The end of %% each line is marked with a \\\\. Lines will be numbered unless the \\\\ %% is preceded by a \\nonumber command. %% Alignment points are marked by ampersands (&). There should be two %% ampersands (&) per line.   %% This section contains more display math examples, including unnumbered %% equations (displaymath environment). The last paragraph includes some %% examples of in-line math featuring a couple of the AASTeX symbol macros.       %% The displaymath environment will produce the same sort of equation as %% the equation environment, except that the equation will not be numbered %% by LaTeX.   %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the"
    },
    "1005/1005.0639_arXiv.txt": {
        "abstract": "We recently predicted the existence of  random primordial magnetic fields (RPMF) in the form of randomly oriented cells with dipole-like structure with a cell size $L_0$ and an average magnetic field $B_0$. Here we investigate models for  primordial magnetic field with a similar web-like structure, and other geometries,  differing perhaps in $L_0$ and $B_0$. The effect of RPMF on the formation of the first galaxies is investigated. The filtering mass, $M_F$, is the halo mass below which baryon accretion is severely depressed. We show that these RPMF could influence the formation of galaxies by altering the filtering mass and the baryon gas fraction of a halo, $f_g$. The effect is particularly strong in small galaxies. We find, for example, for a comoving $B_0=0.1\\muG$, and a reionization epoch that starts at $z_s=11$ and ends at $z_e=8$, for $L_0=100\\,\\text{pc}$ at $z=12$, the $f_g$ becomes severely depressed for $M<10^7\\msun$, whereas for $B_0=0$ the $f_g$ becomes severely depressed only for much smaller masses, $M<10^5\\msun$. We suggest that the observation of $M_F$ and $f_g$ at high redshifts can give information on the intensity and structure of primordial magnetic fields. ",
        "introduction": "Understanding the details of galaxy formation remains an important challenge in cosmology. As shown by numerical calculations, the first generation of galaxies should have formed at very high redshifts inside collapsing halos, starting at $z \\sim 65$, corresponding to high peaks of the primordial dark matter (DM) density field  \\citep{NNB}. Cosmic Microwave Background (CMB) radiation observations  suggest that reionization began at high redshifts. This means that a high abundance of luminous objects must have existed at that time, since these first luminous objects are expected to have heated and reionized their surroundings \\citep{rev,WL03,HH03,Cen}.  The formation of a luminous object inside a halo necessarily  requires the existence of baryonic gas there. Even in halos that are too small for cooling via atomic hydrogen, the gas content can have substantial, and observable, astrophysical effects. In addition to the possibility of hosting astrophysical sources, such as stars, small halos may produce a 21-cm signature \\citep{Kuhlen,Shapiro06,NB08, Furlanetto06}, and can block ionizing radiation and produce an overall delay in the global progress of reionization \\citep{bl02, iliev2, iss05, mcquinn07}.  The evolution of the halo gas fraction at various epochs of the Universe is of prime importance, particularly in the early Universe. We  evaluate here the possible influence of a primordial magnetic field on the halo gas fraction.  As noted by  \\citet{Gnedin2000a, Gnedin1997}, both in the linear and non-linear regimes, the accretion of gas into DM halos is suppressed below a characteristic mass scale called the filtering mass, $M_F$. This mass scale coincides with the Jeans mass, $M_J$, if the latter does not vary in time. Otherwise, $M_F$ is a time average of $M_J$. Thus, an increase in the ambient pressure  in the past, causes an increase in $M_J$ and suppresses the accretion of baryons into DM halos in a cumulative fashion, producing an increase in $M_F$.  Until now, studies focused on the UV heating of the neutral interstellar gas as the main source of pressure, for determining  the  filtering mass. These results  are widely used in many semi-analytic models (e.g.  \\citealt{Maccio2009}), particularly those designed to study the properties of small galaxies (due to the high redshift character of the UV heating).  In this paper we add the effect of a possible random primordial magnetic field as another important source of ambient pressure. The magnetic field contributes to pressure support, which changes the Jeans mass and, consequently, the filtering mass and the quantity of gas that is accreted by DM halos.  The paper is organized as follows. In section \\ref{sec:magnetic} we make a short review on the possible origins of primordial magnetic fields, in section \\ref{sec:filtering} we analyze the effect of primordial magnetic fields on the Jeans and filtering  masses and in section \\ref{sec:gas} we calculate effects on the baryon mass fraction. In section \\ref{sec:conc} we give our conclusions. ",
        "conclusions": "\\label{sec:conc}  We modified the Jeans mass in order to take into account  the presence of  random primordial magnetic fields (RPMF) in the form of randomly oriented cells with dipole and ring-like structures. From this modified Jeans mass, we obtained the filtering mass and the baryonic gas fraction of a dark matter halo. We showed that, depending on the magnetogenesis  model, which determines $B_0$ and $L_0$,  both the Jeans mass and the baryonic gas fraction can change by orders of magnitude. We found, for example, for a comoving $B_0=0.1\\muG$, and a reionization epoch that starts at $z_s=11$ and ends at $z_e=8$, for $L_0=100\\,\\text{pc}$ at $z=12$, the $f_g$ becomes severely depressed for $M<10^7\\msun$, whereas for $B_0=0$ the $f_g$ becomes severely depressed only for much smaller masses, $M<10^5\\msun$.  Since it is very difficult to make observations of intergalactic magnetic fields at high redshifts, and the constraints imposed by CMB measurements are not very restrictive, we suggest the possibility to add new constraints on a family of models for the primordial magnetic field, by following the  redshift evolution of the filtering mass of galaxies.  We also calculated the modified baryonic gas fraction that can also be used as an indirect observable to help us to understand the origin and structure of cosmic magnetic fields."
    },
    "1005/1005.5700_arXiv.txt": {
        "abstract": "To understand the basic mechanism of the formation of magnetic flux concentrations, we determine by direct numerical simulations the turbulence contributions to the mean magnetic pressure in a strongly stratified isothermal layer with large plasma beta, where a weak uniform horizontal mean magnetic field is applied. The negative contribution of turbulence to the effective mean magnetic pressure is determined for strongly stratified forced turbulence over a range of values of magnetic Reynolds and Prandtl numbers. Small-scale dynamo action is shown to reduce the negative effect of turbulence on the effective mean magnetic pressure. However, the turbulence coefficients describing the negative effective magnetic pressure phenomenon are found to be converged for magnetic Reynolds numbers between 60 and 600, which is the largest value considered here. In all these models the {\\it turbulent intensity} is arranged to be nearly independent of height, so the kinetic energy density decreases with height due to the decrease in density. In a second series of numerical experiments, the turbulent intensity increases with height such that the {\\it turbulent kinetic energy density} is nearly independent of height. Turbulent magnetic diffusivity and turbulent pumping velocity are determined with the test-field method for both cases. The vertical profile of the turbulent magnetic diffusivity is found to agree with what is expected based on simple mixing length expressions. Turbulent pumping is shown to be down the gradient of turbulent magnetic diffusivity, but it is twice as large as expected. Corresponding numerical mean-field models are used to show that a large-scale instability can occur in both cases, provided the degree of scale separation is large enough and hence the turbulent magnetic diffusivity small enough. ",
        "introduction": "In a stratified layer, magnetic fields do not normally stay in equilibrium but tend to become buoyantly unstable \\citep[e.g.][]{New61,Par66,Par79a,Gil70a,Gil70b,HP88,Catt88,Wis00,Isobe05,Kers07}, see also reviews by \\cite{Hugh07} and \\cite{Tob07}. This mechanism related to magnetic buoyancy, is generally invoked in order to understand magnetic flux emergence at the solar surface \\citep[e.g.][]{HAGM09}. The mechanism does not explicitly rely upon the existence of turbulence, except that the origin of the Sun's magnetic field is generally believed to be turbulent in nature; see \\cite{Solanki} for a recent review.  Turbulent dynamos work under a variety of circumstances and are able to produce weakly nonuniform large-scale magnetic fields (see Brandenburg \\& Subramanian 2005 for a review). At first glance this generation process is counter-intuitive, because it works against the well-known concept of turbulent mixing \\citep{Tay21,Pra25}. However, it is now well established that turbulence can also cause non-diffusive effects. In addition to the well-known $\\alpha$ effect that is generally believed to be responsible for the Sun's large-scale field \\citep[][]{Mof78,Par79b,KR80}, there is also the $\\Lambda$ effect that is responsible for driving the differential rotation of the Sun \\citep{Rue80,Rue89,RH04}. Yet another import effect is turbulent pumping or $\\gamma$ effect \\citep{Rad69}, which corresponds to the advection of mean magnetic field that is not associated with any material motion. The $\\gamma$ effect appears, for example, in nonuniform turbulence and transports mean magnetic field down the gradient of turbulent intensity, which is usually downward in turbulent convection. However, this effect can also be modified by the mean magnetic field itself \\citep{Kitch94,RK06}, which can then correspond to a mean-field buoyancy effect.  When invoking the concept of magnetic buoyancy, one must ask what the effect of turbulence is in this context. The turbulent pressure associated with the convective fluid motions and magnetic fluctuations is certainly not negligible and reacts sensitively to changes in the background magnetic field. The main reason for this is that the kinetic energy density in isotropic turbulence contributes to the total turbulent dynamic pressure twice as much as turbulent magnetic energy density \\citep[][hereafter referred to as RK07]{KRR89,KRR90,RK07}: \\EQ \\Pturb={\\textstyle{1\\over3}}\\overline{\\rho \\uu^2} +{\\textstyle{1\\over 6}}\\overline{\\bb^2}/\\mu_0. \\EN Here, $\\Pturb$ is the total turbulent dynamic pressure caused by velocity and magnetic fluctuations, $\\uu$ and $\\bb$, respectively, $\\mu_0$ is the vacuum permeability, $\\rho$ is the fluid density, and overbars indicate ensemble averaging. On the other hand, any rise in local turbulent magnetic energy density must be accompanied by an equal and opposite change of turbulent kinetic energy density in order to obey approximate energy conservation, i.e.\\ \\EQ \\half\\overline{\\rho\\uu^2}+\\half\\overline{\\bb^2}/\\mu_0\\equiv E_{\\rm tot}\\approx\\const. \\EN Direct numerical simulations in open systems with boundaries \\citep[][hereafter referred to as BKR]{BKR10} show that when the mean magnetic field $\\meanB$ is much smaller than the equipartition field strength, $\\Beq$, the total energy is conserved, while when $\\meanB \\leq \\Beq$, $E_{\\rm tot}$ decreases slightly with increasing mean field, and it varies certainly less than either $\\overline{\\rho\\uu^2}$ or $\\overline{\\bb^2}$; see Fig.~1 of BKR. This clearly implies that, upon generation of magnetic fluctuations, the total turbulent dynamic pressure shows a reversed (destabilizing) feedback \\citep{KRR90}, i.e.\\ \\EQ \\Pturb=-\\onesixth\\overline{\\bb^2}/\\mu_0+2E_{\\rm tot}/3, \\label{phenomenology} \\EN so both an increase of $\\overline{\\bb^2}$, as well as an increase of the imposed field, which decreases $E_{\\rm tot}$, tend to lower the value of $\\Pturb$. For strongly anisotropic turbulence, \\Eq{phenomenology} is also valid except for the change of the $1/6$ factor into $1/2$ (RK07). This phenomenology was supported by analytical studies using the spectral $\\tau$ relaxation approximation \\citep{KRR90} and the renormalization approach \\citep{KR94} and led to the realization that the effective mean magnetic pressure force (the sum of turbulent and non-turbulent contributions) is reduced and can be reversed for certain mean magnetic field strengths. Under certain conditions (e.g.\\ strong density stratification), this can cause a magnetic buoyancy instability via perturbations of a uniform mean magnetic field in stratified turbulence \\citep{KMR93,KMR96}. Later, when considering the effect of turbulent convection on the mean Lorentz force, RK07 suggested that magnetic flux concentrations in the Sun such as active regions and even sunspots might be formed by this reversed feedback effect.  Most of the numerical simulations on magnetic flux emergence \\citep[e.g.][]{Stein01,Schus06,Mart08,Remp09} have been done using initial conditions with an already existing strongly inhomogeneous large-scale magnetic field. Recent simulations by BKR and \\cite{KKWM10,Kapy11} study the formation of large-scale magnetic structures from an initially uniform large-scale magnetic field. In particular, Large-Eddy Simulations of solar magneto-convection by \\cite{KKWM10} give indications that the spontaneous formation of long-lived magnetic flux concentrations from an initial vertical uniform magnetic field might be possible, although the underlying mechanism in their simulations still remains to be clarified. A similar type of magnetic flux concentration with vertical imposed field has been seen in convection simulations at large aspect ratios by \\cite{Tao98}, which show a segregation into magnetized and weakly magnetized regions. One of the differences compared with BKR is the vertical orientation of the imposed magnetic field in turbulent convection. In forced and convection-driven turbulence simulations of BKR and \\cite{Kapy11}, respectively, the imposed magnetic field was a horizontal one. Other possibilities for causing flux concentrations include turbulent thermal collapse, whereby the magnetic field suppresses the convective energy flux, leading to local cooling, and thus to contraction and further enhancement of magnetic flux \\citep{KM00}. By considering an isothermal equation of state with isothermal stratification, we will exclude this possibility in our present work, allowing thus a more definitive identification of the effect of density-stratified turbulence on the mean Lorentz force.  Meanwhile, direct numerical simulations of forced turbulence with an imposed horizontal magnetic field have demonstrated conclusively that in a simulation with an isothermal equation of state and an isothermal density stratification, spontaneous formation of magnetic structures \\citep{BKKMR11,KBKMR} does indeed occur. Those simulations used a scale separation ratio of 1:15 and 1:30, i.e., the computational domain must be big enough to encompass at least 15 (or even 30) turbulent eddies in one coordinate direction. It does then become computationally expensive to achieve large Reynolds numbers, because their value is based on the size of the energy-carrying turbulent eddies rather than the size of the computational domain. In the present paper we restrict ourselves to a scale separation ratio of 1:5 and are thereby able to demonstrate convergence of the turbulence coefficients describing the negative effective magnetic pressure phenomenon for magnetic Reynolds numbers between 60 and 600.  The basic phenomenon of magnetic flux concentration by the effect of turbulence on the mean Lorentz force has been studied by BKR based on numerical solutions of the {\\it mean-field} momentum and induction equations. They demonstrated the existence of a linear instability for sufficiently strong stratification. This instability was followed by nonlinear saturation at near-equipartition strengths. Using direct numerical simulations (DNS) of forced turbulence, BKR also verified the validity of the phenomenology highlighted by \\Eq{phenomenology}. However, their DNS ignored the effects of stratification which would lead to additional effects such as turbulent pumping that might oppose the instability.  Extending the DNS of BKR to the case with stratification is therefore one of the main goals of the present paper. This will allow us to make a meaningful comparison between DNS in a stratified fluid with mean-field modeling. We are now also able to present data for cases in which small-scale dynamo action is possible. This requires that the magnetic Reynolds number is large enough. As alluded to above, it is then advantageous to choose a scale separation ratio that is not too extreme. While structure formation by the negative effective magnetic pressure phenomenon becomes impossible for small scale separation ratios, there is then also the advantage that the analysis becomes more straightforward in that horizontal and time averages can be employed. This would become problematic in the presence of structures that would break the assumptions of stationarity and homogeneity in the horizontal direction.  There are two other possible caveats that may result from the simplification of using an isothermal equation of state with isothermal stratification. First, the effects of convection and a convectively unstable stratification on the mean Lorentz force are ignored. Fortunately, those turn out to be weak, as shown in a separate paper by \\cite{Kapy11}. Second, owing to the spatial uniformity of the forcing function, $\\overline{\\uu^2}$ is nearly uniform, so the effects of turbulent pumping down the gradient of $\\overline{\\uu^2}$ are ignored. We refer to these as models of type U. Furthermore, because of strong stratification of $\\rho$, the equipartition field strength $\\Beq=(\\mu_0\\overline{\\rho\\uu^2})^{1/2}$ also varies. This has the advantage that a single simulation with imposed field $B_0$ spans a large range in the relevant control parameter $B_0/\\Beq$. However, in view of applications to turbulence in stellar convection this is unrealistic, because there $\\overline{\\rho\\uu^3}$ is nearly independent of height, so $\\Beq$ increases with depth only like $\\rho^{1/6}$. For this reason we also study models in which $\\Beq$ is nearly constant. We refer to these as models of type B. For this purpose we determine first the relevant turbulent pumping velocity which is then used in a suitably adapted mean-field model.  We begin by discussing first the determination of turbulent transport coefficients in \\Sec{DNSmodel}, present our results in \\Sec{Results}, focusing especially on models of type U, turn then in \\Sec{ComparisonUB} to models of type B, compare the results at the level of mean-field models, and finish with a discussion of the main differences between the magnetic buoyancy instability and the negative effective magnetic pressure instability (NEMPI), before concluding in \\Sec{Concl}. ",
        "conclusions": "\\label{Concl}  Our DNS have shown that for an isothermal atmosphere with strong density stratification the total turbulent pressure is decreased due to the generation of magnetic fluctuations by the tangling of an imposed horizontal mean magnetic field by the velocity fluctuations. This phenomenon strongly affects the mean Lorentz force so that the effective mean magnetic pressure becomes negative. For our numerical model with approximately uniform turbulent rms velocity, the ratio of imposed to equipartition field strength changes with height, because the density decreases with height, while the imposed field is constant. This allows us to determine the full functional form of the effective mean magnetic pressure as a function of normalized field strength for a single run.  The form of the dependence of $\\Peff(\\meanB/\\Beq)$ is similar to that found for simulations under rather different conditions (with or without stratification, with or without convection, etc). This dependence is found to be similar to that obtained earlier using both analytic theory (RK07) and direct numerical simulations (BKR), and the results are robust when changing the strength of the imposed field.  In simulations where the turbulent velocity is nearly independent of height, the reduction of magnetic fluctuations occurs in the upper layers where the equipartition field strength decreases with height (models of type U). In models of type B, where the equipartition field strength is nearly constant in height, the magnetic fluctuations are found to be slightly stronger in the upper parts.  In view of astrophysical applications, it is encouraging that $\\qpz$ and $\\betap$ seem to approach an asymptotic regime for $\\Rm>60$. While it remains important to confirm this result, a number of other aspects need to be clarified. Firstly, the issue of finite scale separation is important, i.e., the larger the wavenumber of structures in the mean field relative to $\\kf$, the less efficient the negative effective magnetic pressure will be. This needs to be quantified. For example in the work of \\cite{BKKMR11}, where we had a scale separation ratio of 1:15, magnetic structures were best seen after averaging along the direction of the mean field. On the other hand, with a scale separation ratio of 1:30, structures where quite pronounced already without averaging \\citep{KBKMR}. In the Sun, the scale separation ratio is very large in the horizontal direction, but in the vertical direction the system is extremely inhomogeneous and the vertical pressure scale height increases rapidly with depth. The significance of such effects on the negative effective magnetic pressure effect remains still quite unclear.  Another aspect concerns the limitations imposed by the use of an isothermal equation of state. In the context of regular (non-turbulent) magnetic buoyancy, the system is known to be more unstable to the buoyancy instability when the fluctuations evolve isothermally \\citep{Ach78,HW95,MC03}. However, in the context of the negative effective magnetic pressure instability it is not clear in which direction this effect would work. The only simulations where the equation for specific entropy was taken into account are the simulations of \\cite{Kapy11}, who also considered an unstably stratified atmosphere. In their case, the negative effective magnetic pressure was found to be greatly enhanced (deeper minimum of $\\Peff$ and larger values of $\\Bmin$). Again, this is a subject that deserves serious attention.  The fact that the values of $\\qpz$ and $\\betap$ appear to be converged in the range $60<\\Rm<600$ is significant, because this is also the regime in which small-scale dynamo action occurs. Small-scale dynamo action suppresses the magnetic pressure effect, which is the reason for the drop of $\\qpz$ between $\\Rm$ of 40 and 60, but for larger $\\Rm$, the values of $\\qpz$ seem roughly unchanged.  In the present paper we have discussed applications mainly to the Sun. However, any hydromagnetic turbulence with strong stratification and large plasma beta should be subject to the negative effective magnetic pressure phenomenon. Another relevant example might be accretion disks. Quasi-periodic oscillations and other light curve variations from accretion disks have long been suspected to be caused by some kind of structures in these disks. Hydrodynamic vortices would be one possibility \\citep{Abramo}, which could constitute long-lived structures \\citep{Barge,Johansen,Lyra}. However, in view of the present results, structures caused by the negative effective magnetic pressure instability might indeed be another candidate."
    },
    "1005/1005.0580_arXiv.txt": {
        "abstract": "Hydrogen recombination is one of the most important atomic processes in many astrophysical objects such as Type II supernova (SN~II) atmospheres, the high redshift universe during the cosmological recombination era, and H II regions in the  interstellar medium.  Accurate predictions of the ionization fraction can be quite different from those given by a simple solution if one takes into account many angular momentum sub-states, non-resonant processes, and calculates the rates of all atomic processes from the solution of the radiative transfer equation instead of using a Planck function under the assumption of thermal equilibrium.  We use the general purpose model atmosphere code \\texttt{PHOENIX} 1D to compare how the fundamental probabilities such as the photo-ionization probability, the escape probability, and the collisional de-excitation probability are affected by the presence of other metals in the environment, multiple angular momentum sub-states, and non-resonant processes. Our comparisons are based on a model of SN 1999em, a SNe Type II, 20 days after its explosion. ",
        "introduction": "Hydrogen recombination plays a very important role in many astrophysical phenomena such as Type II supernovae, the interstellar medium, and cosmic recombination.   Time-dependent recombination has been studied for the case of cosmological recombination \\citep{Zeldovich,Peebles} and for supernovae \\citep*{Chugai,dessart08,soma09}. In supernovae this occurs when  the hydrogen recombination time-scale becomes comparable to the age of the supernova. This effect is found to be dominant in the early epochs of the supernova's evolution \\citep{soma09}. At later times, however, effects due to multi-level atom effects (the importance of having many angular momentum sub-states) become more important than time-dependent phenomena. \\citet{soma09} discussed how the effective recombination timescale can be different based on time-dependent rate equations using different hydrogen atom models. The primary goal of \\citet{soma09} was to determine the epoch in the lifetime of a supernova (during the photospheric phase) where time-dependence in the rate equations is most important. In doing so we found  that at later times model atoms with significantly more energy levels (that is additional angular-momentum sub-states) have a strong effect in determining the effective recombination time-scale.  This issue is also important for applications other than supernovae. In fact,  considering more complete atomic models is important to correctly estimate the electron density and recover  subtle features in the spectra. Here, we study multi-level atomic systems for hydrogen alone using a non-LTE treatment which could alter the hydrogen ionization fraction and therefore produce a different temperature structure.  Cosmological recombination codes such as RECFAST \\citep{SSS99,SSS00,WMS08} and RICO \\citep{FCRW09} that deal with cosmological recombination and solve for the free electron fraction as a function of redshift do an excellent job. Nevertheless there are assumptions about Saha equilibrium between higher angular momentum states. \\citet{grin_hirata10} have presented their recombination code RECSPARSE which is a highly detailed calculation but still uses the Sobolev approximation for line calculations. Recently the cosmological recombination problem has been revisted by several authors, paying special attention to possibly neglected atomic physics effects. The processes studied have included two-photon decays in H~I and He~I  \\citep{DG05,WS07,CS08b,Hirata08,KI08,HS08}, H~I continuum absorption of the He I $\\lambda 584$ line photons \\citep{SH08a,SH08b,RCS08,KIV07}, stimulated two-photon decays and two-photon absorption \\citep{Hirata08,CS06,KI08}, Raman scattering \\citep{Hirata08,HS08}, a more complete set of angular momentum substates \\citep{grin_hirata10}, and forbidden transitions in He~I \\citep{WS07,SH08a,SH08b,RCS08}. \\citet{HF09} studied the effects of partial frequency redistribution, line blanketing, and time-dependence in Ly--$\\alpha$, but even then made use of the Sobolev approximation. Here we resolve all line profiles in the co-moving frame and accurately treat line blanketing. We  assume complete redistribution for the line profiles.   We focus primarily on estimating the photo-ionization rate, the escape probability, and the collisional de-excitation rate. We show how these rates depend on the  hydrogen model  atom used and we also study the effects of the presence of metals on the results. These different rates ultimately lead to variations in the hydrogen ionization fraction. We also examine when  the 2$\\gamma$ process plays the most significant role in controlling hydrogen recombination.  In \\S~\\ref{sec:framework}, we outline the theoretical framework relating the photo-ionization rate, collisional de-excitation rate, and escape probability to the effective hydrogen recombination time. In \\S~\\ref{sec:modelatoms}, we describe our approach and different test systems that we study.  In \\S~\\ref{sec:results} we present our results. In \\S~\\ref{sec:discuss}, we state our conclusions and describe a possible framework for future work. ",
        "conclusions": "\\label{sec:discuss}  The primary motivation for this paper was to investigate how the multi-level structure of the hydrogen model atom  affects the hydrogen ionization fraction as well as the photo-ionization, escape, and collisional de-excitation rates. We also investigated if the the metals in the environment or if the  2$\\gamma$ process plays a significant role in the recombination process of the system. We find that a multi-level hydrogen atom structure, or in other words  a more complete set of energy levels to be important in determining the ionization profile of the system. Recently, \\citet{grin_hirata10} studied the effects of including a more complete set of angular momentum substates in the study of cosmological recombination and also found that including more substates to be important. The need to have a more complete set of bound states in the hydrogen atom is found to be more important in a metal-rich environment. This is because of the fact that recombination is more effective in a multi-level framework. When there is a source of additional free electrons (such as metals), the suppression in $f_{H}$ is larger than in pure hydrogen models. This is merely due to the fact that the larger free electron density drives recombination, thus it is very important to use multi-level hydrogen atoms especially in realistic environments with solar compositions.    The enhancement of recombination is indicated by the increasing escape probability for the higher bound states (for both Models B and D). This is also reinforced by the almost constant photo-ionization rate over different energy bound states.   The importance of the 2$\\gamma$ process is seen in the 4-level pure hydrogen case (Model A) as only a small reduction in the free electron density. In Model B, the reduction in the free electron density due to the inclusion of the 2$\\gamma$ process is much larger. There was not any significant effect due to the 2$\\gamma$ process in the solar composition Models. The larger change in Model B (at lower $\\tau_{std}$) can be explained from Figure \\ref{bi_1}. At low optical depths the relative population of the ground state to the first excited state is displayed. For the metal-free Models (A and B) the ratio $\\frac{n_{1}}{n_{2}}$ is not very high and about 4.0 for Model B. Thus in Model B, the first excited state population is not significantly lower than the ground state population. The relative first excited state population is much lower in Models C and D (where the ratio is around $10^{4}-10^{5}$ at similar $\\tau_{std}$).  The high level population of the first excited state (for Model B) increases the importance of the 2$\\gamma$ process, due to the higher number of available electrons that can recombine into the ground state and the upward transition probability for this non-resonant process being low, thus the recombination obtained is more effective. We find that it is essential to incorporate both multi-level atoms and the 2$\\gamma$ process. Our primary purpose was to do a simple calculation to emphasize the important factors in determining the recombination of hydrogen at typical free electron densities found in Type II supernova atmospheres. Our results show that it is important to incorporate 2$\\gamma$ transitions and multiple angular momentum sub-states at low optical depths where the free electron density is small (typically $\\sim 10^{5}$~cm$^{-3}$) and is also typical of the free electron density during the epoch of cosmological recombination \\citep{Peebles}. In H II regions the electron density is much lower making the hydrogen recombination time scale rise. Therefore accurate treatment of hydrogen recombination is important in all of these scenarios. In the future we plan to investigate hydrogen recombination in the context of cosmological recombination epoch."
    },
    "1005/1005.1779_arXiv.txt": {
        "abstract": "Multistate dark matter (DM) models with small mass splittings and couplings to light hidden sector bosons have been proposed as an explanation for the PAMELA/Fermi/H.E.S.S.\\ high-energy lepton excesses.  We investigate this proposal over a wide range of DM density profiles, in the framework of concrete models with doublet or triplet  dark matter and a hidden SU(2) gauge sector that mixes with standard model hypercharge.   The gauge coupling is bounded from below by the DM relic density, and the Sommerfeld enhancement factor is explicitly computable for given values of the DM and gauge boson masses $M$, $\\mu$ and the (largest) dark matter mass splitting $\\delta M_{12}$.   Sommerfeld enhancement is stronger at the galactic center than near the Sun because of the radial dependence of the DM velocity profile, which strengthens the inverse Compton (IC) gamma ray constraints relative to usual assumptions.  We find that the PAMELA/Fermi/H.E.S.S.\\ lepton excesses are marginally compatible with the model predictions, and with CMB and Fermi gamma ray constraints, for $M\\cong 800$ GeV, $\\mu\\lsim 200$ MeV, and a dark matter profile with noncuspy Einasto parameters $\\alpha\\gsim 0.20$, $r_s \\sim 30$ kpc. We also find that the annihilating DM must provide only  a subdominant ($\\lsim 0.4$) component of the total DM mass density, since otherwise the boost factor due to Sommerfeld enhancement is too large. ",
        "introduction": "Dark matter (DM) annihilating in the galaxy has been proposed as the explanation for a number of current experimental anomalies.   The Payload for Antimatter-Matter Exploration and Light-nuclei Astrophysics (PAMELA) experiment has observed an excess in the positron fraction in the energy range $10-100$ GeV \\cite{pamela}, and the Fermi Large Area Telescope (LAT) \\cite{fermi} observes an excess in $e^++e^-$ at energies up to 1 TeV, compared to the predictions of conventional diffusive models. The High Energy Stereoscopic System (H.E.S.S.) \\cite{hess}, too, observes a steepening in the spectrum around a few TeV which is in agreement with the Fermi-LAT excess. On the other hand, PAMELA does not observe excess antiprotons.  Therefore  TeV-mass particles that can annihilate preferentially into leptons are the preferred DM candidates for explaining these anomalies \\cite{CKRS,donato,cholis,berg,meade}.   Theories of leptophilic annihilating DM have come under increasing pressure from a number of complementary constraints, due to the radio and gamma ray (GR) emissions that should be produced by the annihilation process itself (`prompt' gamma rays) \\cite{prompt1,prompt2,prompt3,prompt4} or by the high energy $e^\\pm$, via synchrotron radiation in the galactic magnetic field or via inverse Compton scattering on cosmic microwave background (CMB) photons and the galactic radiation field (starlight) \\cite{CP,cuoco,Regis:2009md,Cholis:2009gv,papucci,CPS,zaharijas,hutsi}.  Such radiation could also partially reionize the universe after recombination, leading to distortions in the CMB \\cite{galli,SPF,hutsiCMB,CIP,kanzakiCMB}. Finally, relevant constraints are imposed by extragalactic gamma rays, produced by the annihilations in forming DM halos \\cite{Profumo,Belikov:2009cx,hutsiCMB}. Many models of annihilating DM are seemingly ruled out by overproducing such radiation. In particular, inverse Compton gamma rays produced in the galaxy would contribute to the GR spectrum observed by Fermi-LAT.  A precise assesment of the strength of Fermi inverse Compton (IC) constraints depends however upon the details of the assumed DM density profile $\\rho(r)$. The strongest constraints come from close to the Galactic Center (GC); therefore they can be satisfied if $\\rho(r)$ is not strongly peaked as $r\\to 0$, where $r$ is the distance from the GC. Usually such dependence is illustrated by choosing a few different parametrizations of the profile with varying cuspiness ({\\it e.g.,} the peaked Navarro-Frenk-White profile \\cite{NFW} as opposed to the cored isothermal \\cite{isothermal} or Burkert profiles \\cite{burkert}). $N$-body simulations and observations of our own galaxy suggest a reasonable range of shapes for $\\rho(r)$, and it would be interesting to explore the dependence of the constraints on the characteristics of $\\rho(r)$  within such a range.  One purpose of the present work is to fill this gap by quantifying the strength of the IC constraint for a wide range of DM density profiles.   We limit ourselves to the Einasto parametrization \\cite{Graham:2005xx,Navarro2008}, which we find to be useful for this purpose since it  allows for continuous variation in the cuspiness and it has been used to fit the results of high-resolution $N$-body simulations.  \\smallskip  The constraints also depend upon which leptons, and how many of them, are assumed to be produced in the DM annihilations.  For example, models that produce only $e^+e^-$ are more strongly constrained, for a given DM mass $M$, than those that produce two $e^\\pm$ pairs, since the spectrum of the former is harder. This is a relevant distinction since some of the  models that are most natural from the  particle physics point of view  necessarily produce the $4e$ (or at least $2e$ plus an invisible hidden sector gauge boson) rather than the $2e$ final state.  Likewise models that produce muons are more strongly constrained than those that produce only electrons.   Moreover models using the gauge kinetic mixing portal produce an admixture of $e^\\pm$, $\\mu^\\pm$ and $\\pi^\\pm$ rather than just a single kind of final state. Failing to recognize these distinctions, one could be misled into thinking that the general class of models is excluded, even though there may exist specific examples that satisfy all the constraints.  Thus another goal of the present work is to explore in detail the viability of a very specific set of annihilating DM models, which are also theoretically motivated \\cite{nima}-\\cite{nonabelian}.  Namely we consider DM $\\chi_i$ transforming under a nonabelian hidden sector gauge symmetry, taking the simplest example of SU(2) and the lowest representations (doublet or triplet) for the DM.  One of the dark gauge bosons $B_i$ is assumed to mix with standard model hypercharge (and thus the photon) through a dimension-5 gauge kinetic mixing operator.   This model has a number of appealing features. The DM annihilates into dark gauge bosons,  $\\chi\\chi\\to BB$.   By assuming the gauge symmetry breaks below the GeV scale, so the $B_i$'s have mass $\\mu\\lsim 1$ GeV, it is assured that only light leptons or $\\pi^\\pm$ are produced by the gauge boson decays, and no antiprotrons that would be in conflict with the PAMELA observations.  The annihilation cross section is naturally enhanced by Sommerfeld effect \\cite{sommerfeld,nima} in this model, allowing it to have the correct relic density in the Early Universe, even though the cross section in the galaxy must be larger by a factor of order 100 than that needed for thermal freeze-out.  A distinctive feature of the model is that small mass splittings $\\delta M\\sim$ 1 MeV between the DM states are generated at one loop.   Such splittings can be important for making the model consistent with direct detection constraints, by causing DM scattering on nucleons to be inelastic and endothermic. This allows the rate of DM interactions in detectors to be  sufficiently suppressed despite having relatively large couplings to nucleons \\cite{nima}.   The mass splittings can moreover lead to a significant increase in the Sommerfeld enhancement factor \\cite{slatyer}, which effect we take into account here.  In addition, they can potentially help to explain the annual modulation observed by DAMA/LIBRA through the inelastic DM mechanism (see for example \\cite{idm}), or the 511 keV excess observed by INTEGRAL \\cite{spi} using the excited DM mechanism \\cite{FW}-\\cite{newccf}, although we do not explore these directions in the present work.  A further refinement we make compared to most previous analyses is to  take into account the dependence of the Sommerfeld enhancement on the position in the Galaxy, an effect that was first pointed out in ref.\\ \\cite{robertson}. The Sommerfeld enhancement depends sensitively on the DM velocity, and both the DM velocity dispersion and escape velocity depend upon $r$. We find that this typically makes the enhancement larger near the galactic center where the gamma rays that are most important for the IC constraint are produced.  The constraint thus becomes harder to satisfy than in models with a spatially constant boost factor, as is usually assumed.  Nevertheless, we find examples of models that are marginally able to produce  leptons consistent with the PAMELA and Fermi observations, and which are also barely consistent with the various gamma ray constraints.  The preferred models have DM with $M\\cong 800$ GeV and gauge boson masses below 200 MeV, so that only electrons and no muons or charged pions are produced in the annihilations.  However we find that the Sommerfeld enhancement factor would always be {\\it too large} to achieve this concordance if the DM was present at the expected relic density.  We are obliged to suppress the rate of annihilations by introducing a factor $1/f$ in the density of the annihilating component.  This can be achieved by increasing the gauge coupling appropriately so that the relic density is reduced.  The minimum reduction $1/f=0.4$ occurs if the DM mass splitting is negligible, and smaller values  $1/f \\cong 0.14-0.2$ are needed if $\\delta M \\sim 1$ MeV.  This comes about because of the $\\delta M$-dependence of the Sommerfeld enhancement factor for multistate DM annihilations \\cite{slatyer}.  In the present work we do not consider astrophysical boost factors due to  increased annihilation rate in substructures of the main halo \\cite{Kuhlen:2008aw}-\\cite{Kamionkowski:2010mi}.  Inclusion of these effects would presumably strengthen the gamma ray constraints on the models,  but on the other hand a fully consistent treatment of substructure effects should simultaneously consider the modification to the  lepton signal from annihilations in subhalos \\cite{Brun:2009aj}.  This could conceivably have a compensating effect that would weaken the constraints due to requiring less $e^\\pm$ production near the galactic center \\cite{subhalo}.  Further investigation of this issue is in progress.   We start by introducing the particle physics models in section \\ref{models}, and recalling the constraint upon the dark SU(2) gauge coupling imposed by the relic density.   In section \\ref{boost_sect} we analyze how large an enhancement factor is needed for the models to explain the anomalous lepton observations, the theoretical computation of the Sommerfeld enhancement, and its phase space average in the galaxy.  Section \\ref{icconst} explains how we implement the inverse Compton $\\gamma$ ray constraint for general DM Einasto profiles, while other relevant constraints are discussed in section \\ref{other}.  Our main results showing which parameter ranges can be compatible with all the observations are presented in section \\ref{results}.  We discuss their implications in section \\ref{discussion}. Three appendices give details for the branching ratio of ground state DM annihilation into leptons (app.\\ \\ref{brfr}), the multistate Sommerfeld enhancement (app.\\ \\ref{bfa}) and the calculation of the DM escape velocity (app.\\ \\ref{escvel}) as a function of radial position. ",
        "conclusions": "\\label{discussion}  We have tried to carefully assess the likelihood that some of the best theoretically motivated models of annihilating multistate dark matter can explain the PAMELA/Fermi/H.E.S.S.\\ lepton excesses, considering important details of the models whose implications could be missed by a more generic, model-independent investigation.   As well, we take into account simulations and observations of the DM halo properties, which have a crucial impact on the viability of the scenario.  The main challenge to the models is the difficulty of satisfying constraints from Fermi due to the production of inverse Compton (IC) $\\gamma$ rays produced by the high-energy leptons scattering on galactic radiation.  Among the most important model-dependent features, which affect the ease of  satisfying the constraints, are that each annihilation produces two dark gauge bosons, hence four leptons, and these are all $e^\\pm$ as opposed to heavier charged particles if the gauge boson masses are lighter than $\\sim 210$ MeV.  The IC constraints for this case are less severe than for models that produce only two leptons, or ones that produce $\\mu^\\pm$ in addition to $e^\\pm$.   Moreover the gauge coupling is bounded from below by the requirement that the relic density not exceed what is observed.  We find that the boost factor from Sommerfeld enhancement is too large to satisfy the constraints, unless the density of the annihilating DM is suppressed by a factor $1/f < 0.4$; hence one would need an additional component of nonleptophilic DM to make up the rest.  Using recent results for the Sommerfeld enhancement factor of multistate DM \\cite{slatyer}, this problem is exacerbated when the mass splitting is nonzero; one needs $1/f \\sim 0.14-0.2$ if $\\delta M \\sim 1$ MeV.  A main result is that the Einasto parameters which determine the shape of the DM halo must take large values $\\alpha \\gsim 0.20$ and $r_s \\sim 30$ kpc in order to barely satisfy the IC constraint and have marginal consistency with the observed lepton excesses.   In the context of pure dark matter simulations, such a large value of $\\alpha$ is above the norm for a galaxy of the size of the Milky Way.   Ref.\\  \\cite{gao} finds a correlation between $\\alpha$ and the virial mass of galaxies $M_{\\rm vir}$, $\\alpha = 0.155 + 0.0095\\, \\nu^2$, where  $\\nu = \\delta_{\\rm crit}(z)/\\sigma(M_{\\rm vir})$ are functions arising in the Press-Schechter formalism \\cite{press-schechter} such that  $\\delta_{\\rm crit} = 1.686$ at redshift $z=0$, and $\\sigma > 2$  \\cite{loeb} for a galaxy such as ours with $M_{\\rm vir} = 10^{12} M_\\odot$ \\cite{xue}.  Thus $\\nu < 0.843$, which implies  $\\alpha < 0.162$.  However this is only true for the mean value, and a fluctuation as large as $\\alpha=0.20$ for a galaxy like ours is possible \\cite{navarro-com}.  On the other hand,  recent results incorporating baryons indicate that $\\alpha$ is decreased relative to its pure DM value due to adiabatic contraction, which exacerbates the conflict.  \\medskip  We conclude that there is a distinct tension for  annihilating models to consistently account for all the lepton excesses, while it is easier to explain that of PAMELA alone using smaller DM masses $M<400$ GeV.   A possible loophole that we have not considered here is the suggestion that leptons originating from subhalos of our main halo could provide a substantial fraction of the observed  leptonic excess \\cite{subhalo}.  If this is the case, the production  of IC $\\gamma$ rays from the galactic center would be reduced, providing another way to weaken the constraint. How significant this effect could be is under investigation.  \\bigskip  {{\\bf Acknowledgments}. We thank Fang Chen, Ilias Cholis, Julio Navarro, Pasquale Serpico, Tracy Slatyer and Patricia Tissera for helpful discussions and correspondence. We are especially grateful to Alessandro Strumia for detailed information pertaining to ref.\\ \\cite{papucci}. M.C.\\ thanks the EU Marie Curie Research \\& Training network ``UniverseNet\" (MRTN-CT-2006-035863) for support.  J.C.\\ is supported by the Natural Sciences and Engineering Research Council (Canada).}  \\appendix"
    },
    "1005/1005.1854_arXiv.txt": {
        "abstract": "We report the detection of a magnetic field on the Of?p star HD~108. Spectropolarimetric observations conducted in 2007, 2008 and 2009 respectively with NARVAL@TBL and ESPaDOnS@CFHT reveal a clear Zeeman signature in the average Stokes V profile, stable on timescales of days to months and slowly increasing in amplitude on timescales of years.  We speculate that this timescale is the same as that on which H$\\alpha$ emission is varying and is equal to the rotation period of the star. The corresponding longitudinal magnetic field, measured during each of the three seasons, increases slowly from 100 to 150~G, implying that the polar strength of the putatively-dipolar large-scale magnetic field of HD~108 is at least 0.5~kG and most likely of the order of 1--2 kG.  The stellar and wind properties are derived through a quantitative spectroscopic analysis with the code CMFGEN. The effective temperature is difficult to constrain because of the unusually strong He~{\\sc i}~$\\lambda$~ 4471 and He~{\\sc i}~$\\lambda$~5876 lines. Values in the range 33000--37000 K are preferred. A mass loss rate of about $10^{-7}$ \\myr\\ (with a clumping factor f=0.01) and a wind terminal velocity of 2000 \\kms\\ are derived. The wind confinement parameter $\\eta_{\\star}$ is larger than 100, implying that the wind of HD~108 is magnetically confined.  Stochastic short-term variability is observed in the wind-sensitive lines but not in the photospheric lines, excluding the presence of pulsations. Material infall in the confined wind is the most likely origin for lines formed in the inner wind. Wind--clumping also probably causes part of the H$\\alpha$ variability. The projected rotational velocity of HD~108 is lower than 50 \\kms, consistent with the spectroscopic and photometric variation timescales of a few decades. Overall, HD~108 is very similar to the magnetic O star HD~191612 except for an even slower rotation. ",
        "introduction": "\\label{s_intro}  Very little is known about magnetic fields of O stars and very few O stars are yet known as magnetic; preliminary results \\citep[e.g.,][]{donati02,donati06,bouret08,petit08} suggest that most magnetic O stars could be the high-mass equivalent of magnetic chemically peculiar A and B stars (the so-called Ap and Bp stars) with fields of primordial origin, i.e., fossil remnants from the formation stage \\citep{dl09}.  They are also suspected to be progenitors of highly-magnetic neutron stars and magnetars \\citep[e.g.][]{petit08,Ferrario08}, although alternative explanations exist \\citep{davies09}.  Whenever known, rotation rates of magnetic O stars are significantly smaller than non-magnetic O stars of similar spectral types, further strengthening the analogy with the lower-mass Ap and Bp stars.  Evolutionary models of massive stars including magnetic fields also suggest that the internal rotation is strongly affected by the presence of magnetic fields, enforcing in particular mostly solid-body rotation profiles throughout the bulk of the outer radiative zone \\citep{Maeder03,Maeder04}.  Last but not least, magnetic fields are potentially capable of affecting the powerful radiative winds of O stars \\citep{bm97} and, in some extreme cases, of fully dissipating their angular momentum through wind magnetic braking \\citep{ud09}. Lives of magnetic O stars are therefore expected to differ significantly from those of their non-magnetic counterparts.  In this respect, the prototypical cases of \\tori\\ \\citep{donati02} and HD~191612 \\citep{donati06} are very interesting. The magnetic field of \\tori\\ (about 1.1~kG tilted at about 45\\degr\\ to the rotation axis) is strong enough to deflect the wind flows from both magnetic hemispheres towards the magnetic equator where they collide, generating a strong shock, a very hot (coronal-like) post-shock region and a denser cooling disc (more or less confined to the magnetic equatorial plane). Emission in Balmer lines exhibits very strong rotational modulation, the amount of emission being directly related to the geometry of the cooling disc as seen from the Earth: strongest when the disc is seen pole-on and weakest when the disc is seen edge-on \\citep[e.g.,][]{donati02}. The presence of magnetic fields and the very similar \\hal\\ modulation reported for HD~191612 suggests that this star is analogous to \\tori\\ \\citep{donati06}, but more evolved and much more slowly rotating (538~d vs.\\ 15.4~d for \\tori) as a likely result of wind magnetic braking.  Recent (yet unpublished) data (Wade et al, in preparation) fully confirms the initial conclusions and the remarkable similarity between HD~191612 and \\tori.  Exploring further the potential impact of magnetic fields on the life of massive stars requires new magnetic O stars to be identified. From the example of \\tori\\ and HD~191612, those exhibiting periodically varying \\hal\\ emission are likely to be very good candidates. In particular, HD~108 is interesting. It is a member of the Of?p class composed of the Galactic stars HD~191612, HD148937, NGC~1624--2 and CPD -28$\\deg$ 2561, the latter two recently discovered by \\citet{walborn10}. This class is principally defined by emission in the C {\\sc iii} 4647--4650 lines at least as strong as in the neighbouring N {\\sc iii} 4630--4640 lines. Secondary classification criteria include sharp emission or P-Cygni profiles in He {\\sc i} and Balmer lines.  As \\tori\\ and HD~191612, HD~108 shows spectral type variations (from 04 to 08) due to changes in the strength of its He~{\\sc i} lines. They (and the Balmer lines too) fluctuate between pure absorption and P--Cygni profiles. These line morphologies define ``low'' and ``high'' states respectively. The changes are correlated to photometric variations: HD~108 is brighter when the emission components are stronger, again similarly to HD~191612. They seem to occur with a period of a few decades \\citep{naze01,barannikov07} compared to 538 days for HD~191612 \\citep{walborn04,howarth07}. If the analogy with \\tori\\ and HD~191612 holds, HD~108 could thus potentially show up as a rather extreme case of wind magnetic braking. Note that a marginal detection based on a single exposure was reported by \\citet{hubrig08} for the third member of the Galactic Of?p class, HD148937 \\citep[but see][for a discussion of the significance of the detection]{silv09}. Further observations are required to confirm this detection.  To investigate this issue, we initiated a spectropolarimetric monitoring program of HD~108 using NARVAL at the 2~m T\\'elescope Bernard Lyot (TBL) in the French Pyrenees and later on ESPaDOnS at the 3.6~m Canada-France-Hawaii Telescope (CFHT) atop Mauna Kea, Hawaii\\footnote{This effort is part of the MiMeS (Magnetism of Massive Stars) international research program aimed at characterizing magnetic fields in upper main sequence stars.}.  Reports indicate that HD~108 has just passed the epoch of minimum \\hal\\ emission (Naz\\'e et al, in preparation), suggesting that longitudinal magnetic fields are also close to their minimum value and slowly rising again (provided the analogy with \\tori\\ and HD~191612 holds).  We report here the first results after 3 observing seasons (2007 to 2009). In particular, we report the detection of a longitudinal magnetic field of about 100~G at the surface of HD~108, suggesting that the analogy with \\tori\\ and HD~191612 is likely.  We also carry out a detailed spectral analysis to obtain updated constraints on the parameters of HD~108. We finally discuss briefly the implications of our results for understanding of how magnetic fields affect the lives of massive stars. ",
        "conclusions": "\\label{s_conc}  We have presented spectropolarimetric observations of the Of?p star HD~108 conducted with the NARVAL and ESPaDOnS instruments (at TBL and CFHT respectively). One hundred and ten circularly--polarized spectral sequences have been collected between 2007 and 2009. We report the clear detection of a Stokes V Zeeman signature stable on timescales of days to months, but likely slowly increasing in amplitude on timescales of years. We speculate that this timescale is the same as that on which the H$\\alpha$ emission is varying and is equal to the rotation period of the star. The corresponding longitudinal magnetic field is of order of 100--150~G, implying that the polar strength of the putatively-dipolar large-scale magnetic field of HD~108 is at least 0.5~kG and most likely of order of 1--2 kG.  The stellar and wind parameters of HD~108 have been derived through atmosphere modelling with the code CMFGEN. The effective temperature was poorly constrained due to the surprisingly strong He~{\\sc i}~$\\lambda$~4471 and He~{\\sc i}~$\\lambda$~5876 lines. Values of \\teff\\ in the range 33000--37000 K were conservatively derived from the analysis of the remaining He~{\\sc i} and the He~{\\sc ii} lines. A mass loss rate of the order $10^{-7}$ \\myr\\ was derived. The wind is found to be strongly clumped (f=0.01). It is also significantly confined ($\\eta_{\\star} \\geq 100$, possibly up to 800) by the magnetic field. HD~108 shows long-term variability in most He~{\\sc i} and Balmer lines, with profiles changing from pure absorption to P--Cygni on a few decades period. We suggest that this periodicity corresponds to rotational modulation. The (magnetic) equatorial density enhancement implied by the wind confinement would then be seen edge-on at minimum line emission, and pole-on at maximum emission. The timescale for magnetic braking computed with the derived stellar, wind and magnetic properties is consistent with the star's age. A slow rotation is also consistent with the low \\vsini\\ we derive. HD~108 might thus be an even more extreme case of slowly rotating magnetic O star than HD~191612.  Short-term line profile variability is also observed in He~{\\sc i} and wind sensitive lines. Photospheric He~{\\sc ii} lines are very stable. Lines formed just above the photosphere show evidence for infall while the variability of lines formed in the outer wind is more typical of normal hot stars' outflows. We suggest that in the inner wind, we are witnessing the infall of material channeled by the field lines to the magnetic equator and subsequently pulled back to the stellar surface by gravity.  Most of the suggestions discussed in Sect.\\ \\ref{s_disc} and summarized above are still speculative. Future monitoring of HD~108 is obviously needed to confirm the expected correlation between the variation of the longitudinal magnetic field and the long-term spectroscopic variability. If the few decades spectroscopic and photometric modulation corresponds to the rotational period, a complete mapping of the magnetic topology is not possible until a few decades. However, crucial information regarding the field strength and geometry can be gathered in the next years since we have just passed the phase of minimum emission. The field is thus expected to strengthen, making detection easier. Further constraints on the field morphology will help to see 1) if the slow rotation of HD~108 is really a due to magnetic braking and 2) if the short-term variability is caused by dynamical phenomena predicted by current simulations of magnetically--channelled winds."
    },
    "1005/1005.3060_arXiv.txt": {
        "abstract": "{ We use the Science Demonstration Phase data of the \\emph{Herschel} Virgo Cluster Survey to search for dust emission of early-type dwarf galaxies in the central regions of the Virgo Cluster as an alternative way of identifying the interstellar medium. We present the first possible far-infrared detection of cluster early-type dwarf galaxies: \\object{VCC\\,781} and \\object{VCC\\,951} are detected at the $10\\sigma$ level in the SPIRE 250~$\\mu$m image. Both detected galaxies have dust masses of the order of $10^5$~M$_\\odot$ and average dust temperatures $\\approx$ 20K. The detection rate (less than 1\\%) is quite high compared to the 1.7$\\%$ detection rate for H{\\sc{i}} emission, considering that dwarfs in the central regions are more H{\\sc{i}} deficient. We conclude that the removal of interstellar dust from dwarf galaxies resulting from ram pressure stripping, harassment, or tidal effects must be as efficient as the removal of interstellar gas.} ",
        "introduction": "Early-type dwarf galaxies (dEs) are the dominant morphological type in galaxy clusters. They were originally seen as a rather homogeneous population of dwarf galaxies with an old stellar age, no features resembling any recent or ongoing star formation, and no indications of a significant interstellar medium (ISM). This viewpoint has changed radically in the past few years. Deep imaging observations of dEs have revealed a heterogeneous morphology. In particular, the population of dEs can be subdivided into nucleated and non-nucleated subclasses and several papers report evidence of disks, spiral structure, bars or star formation \\citep[e.g.,][]{Jerjen2000, Barazza2002, Geha2003, Graham2003, GrahamGuzman2003, DeRijcke2003, Lisker2006_1, Lisker2006_2}. Kinematical studies demonstrate that dEs are not simple pressure-supported systems: some dEs seem to be rotationally supported \\citep{Simien2002, Pedraz2002, Geha2003, DeRijcke2003, VanZee2004, Toloba2009}, whereas others show evidence of kinematically decoupled cores \\citep{2004A&A...426...53D, 2006A&A...445L..19T}.  Adding to this morphological and kinematical inhomogeneity is the detection of a significant interstellar medium in various dEs. Atomic and molecular gas was discovered for the first time in the Local Group dEs NGC\\,185 and NGC\\,205 using deep VLA and BIMA observations \\citep{1997ApJ...476..127Y}. In the Virgo cluster, \\citet{Conselice2003} estimate a H{\\sc{i}} detection rate of 15\\% for dEs down to a 3$\\sigma$ limit of  $8 \\times 10^{6}$~M$_\\odot$, while \\citet{Alighieri2007} report a 1.7$\\%$ H{\\sc{i}} detection rate with a 3$\\sigma$ detection limit of $3.5 \\times 10^{7}$~M$_\\odot$. Large amounts of cold interstellar matter are unexpected in dEs, as both internal (supernova explosions) and external effects (ram-pressure stripping, galaxy interactions, tidal effects) are thought to be able to expel the gas from the shallow potential on short timescales \\citep{2004ANS...325..122M, VanZee2004, 2005A&A...433..875R, 2008ApJ...674..742B, 2008A&A...489.1015B, Valcke2008}. That the H{\\sc{i}} detected dEs are found preferentially near the edge of the cluster supports the idea that understanding environmental effects is crucial to constraining the evolutionary history of dEs. \\citet{Buyle2005} and \\citet{Bouchard2005} performed H{\\sc{i}} observations in the outskirts of the Fornax cluster and the Sculptor group, respectively, and confirmed the tendency of gas-deficient galaxies to be located near the center of the cluster.  Continuum emission from interstellar dust is a promising alternative way to determine the ISM content of dEs. In particular, the ISM in the immediate surroundings of M87 can be traced by dust emission, as strong radio continuum emission considerably reduces the ALFALFA H{\\sc{i}} detection sensitivity within 1$\\degr$ of M87 \\citep{2007AJ....133.2569G}. To date, the Andromeda satellites NGC\\,205 and NGC\\,185 are the only dEs that have been detected in the far-infrared \\citep{Haas1998, Marleau2006, Marleau2009}. In particular, no dust emission has been detected from cluster dEs. While extinction features in optical images are indicative of dust in at least some cluster dEs \\citep{Ferrarese2006, Lisker2006_1, Lisker2006_2}, the direct detection of dust emission has been hampered by the limited resolution and sensitivity of the previous generation of far-infrared instrumentation. The advent of the \\emph{Herschel} Space Observatory \\citep{Herschel} offers new possibilities for mapping the ISM in dEs. In this paper, we report on our search for far-infrared emission from dEs in the Virgo Cluster, based on the Science Demonstration Phase (SDP) observations of the \\emph{Herschel} Virgo Cluster Survey (HeViCS\\footnote{More details on HeViCS can be found on http://www.hevics.org.}). In Sect.~{\\ref{Observations.sec}}, we describe the observations, data reduction, and sample selection. In Sect.~{\\ref{Detections.sec}}, we present our detections and in Sect.~{\\ref{Stacking.sec}} we discuss the results of a stacking analysis of the non-detections. Sect.~{\\ref{Discussion.sec}} gives our summary.  \\begin{figure*} \\centering \\includegraphics[width=\\textwidth]{14647fg1.pdf} \\caption{ All images are on the same scale. The FWHM of the beam is denoted in the MIPS 24 $\\mu$m and SPIRE 250 $\\mu$m images. g-band contours are overlaid on the SPIRE 250 $\\mu$m image. The dwarfs position is indicated as a star, SDSS background sources as open circles, and the neighbouring source of VCC\\,951 as a cross.} \\label{VCC781.pdf} \\end{figure*} ",
        "conclusions": "\\label{Discussion.sec}  We have reported on our search for far-infrared dust emission in a sample of 239 dEs at the centre of the Virgo Cluster. With a detection limit of about $10^4~M_\\odot$, we have detected two objects at 10$\\sigma$ that have a high probability of being true detections of dEs. These two dwarf galaxies,\\object{VCC\\,781} and \\object{VCC\\,951}, are the first detections of dust emission in dEs, apart from the Andromeda satellites NGC\\,205 and NGC\\,185. Applying a radiative transfer model to determine the equilibrium dust temperature in their interstellar radiation field, we determined temperatures of the order of 20~K in both galaxies. We estimate  dust masses of the order M$_{\\text{d}} = 10^5$~M$_\\odot$ in both dEs.  In addition to independent detections of both atomic and molecular gas in dEs, the detection of dust emission here confirms the hypothesis that at least some dEs are transition objects gradually evolving from late-type to early-type, while falling into the cluster.  The detection of an interstellar dust medium in two dEs is rather unexpected, given the efficiency with which the ISM is removed from dwarf galaxies. \\citet{2008ApJ...674..742B} found that gas removal by ram pressure stripping is extremely efficient in dwarfs galaxies, there being a decrease by $\\sim$2 orders of magnitude in $\\sim$150~Myr of time. In agreement with this rapid morphological transition, only a small number ($10-16 \\%$) of dwarf galaxies exhibiting properties intermediate between star-forming and spheroidal galaxies are present in the Virgo Cluster.  \\citet{Alighieri2007} find an H{\\sc{i}} detection rate in the Virgo cluster of only 1.7\\% to a 3$\\sigma$ mass limit M$_{\\text{H{\\sc{i}}}} = 3.5 \\times 10^7$~M$_\\odot$. Adopting a canonical gas-to-dust ratio of 1000 (which might still be a conservative lower limit as dEs in Virgo are expected to have metallicities significantly lower than observed in our own Milky Way), we recover a $3\\sigma$ dust mass limit of $3.5 \\times 10^4$~M$_\\odot$, which is more or less the $3\\sigma$ dust mass detection limit in the HeViCS SDP field. Our detection rate is also compatible with the 15\\% found by \\citet{Conselice2003}. Their H{\\sc{i}} upper limit corresponds to a 3$\\sigma$ dust mass limit of $8 \\times 10^{3}$~M$_\\odot$, well below the SDP detection limit. When also taking into account the position in the cluster, we consider our detection rate of less than 1\\% surprisingly high. Indeed, H{\\sc{i}} detected dEs are preferentially located in the outskirts of the cluster and may not have been dramatically affected so far by the intracluster medium and interactions with other galaxies \\citep{Conselice2003, Buyle2005}. \\citet{Alighieri2007} also report that the majority of the detected H{\\sc{i}}-sources are located near the edges of the Virgo cluster. If gas and dust are tightly coupled and dust is as easily removed from the galaxy as gas by ram pressure stripping, harassment, or tidal effects, we expect that the dEs in the central regions of the Virgo Cluster will also be the most dust deficient. Clear evidence of efficient dust stripping in the Virgo Cluster is seen for the first time in H{\\sc{i}}-deficient spiral galaxies \\citep{Luca}. Future HeViCS observations will be able to provide an answer to this question -- the HeViCS survey will not only go deeper (each field will be covered 4 times), but also wider (the total area covered will be $\\sim$64 deg$^2$). A comparison of the far-infrared detection rate with the position in the cluster will tell us whether environmental effects have any influence on the detection rate."
    },
    "1005/1005.4706_arXiv.txt": {
        "abstract": "In an effort to understand the factors that govern the transition from low- to high-mass star formation, we identify for the first time a sample of intermediate-mass star-forming regions (IM SFRs) where stars up to---but not exceeding---$\\sim 8$ M$_{\\sun}$ are being produced. We use \\emph{IRAS} colors and \\emph{Spitzer Space Telescope} mid-IR images, in conjunction with millimeter continuum and $^{13}$CO maps, to compile a sample of 50 IM SFRs in the inner Galaxy. These are likely to be precursors to Herbig AeBe stars and their associated clusters of low-mass stars. IM SFRs constitute embedded clusters at an early evolutionary stage akin to compact \\ion{H}{2} regions, but they lack the massive ionizing central star(s). The photodissociation regions that demarcate IM SFRs have typical diameters of $\\sim1$ pc and luminosities of $\\sim 10^4$ L$_{\\sun}$, making them an order of magnitude less luminous than (ultra)compact \\ion{H}{2} regions. IM SFRs coincide with molecular clumps of mass $\\sim 10^3$ M$_{\\sun}$ which, in turn, lie within larger molecular clouds spanning the lower end of the giant molecular cloud mass range, $10^4-10^5$ M$_{\\sun}$. The IR luminosity and associated molecular mass of IM SFRs are correlated, consistent with the known luminosity--mass relationship of compact \\ion{H}{2} regions. Peak mass column densities within IM SFRs are $\\sim$ 0.1--0.5 g~cm$^{-2}$, a factor of several lower than ultra-compact \\ion{H}{2} regions, supporting the proposition that there is a threshold for massive star formation at $\\sim$ 1 g~cm$^{-2}$. ",
        "introduction": "It has become clear that high-mass star formation cannot be viewed simply as a scaled-up version of the low-mass star formation paradigm \\citep{zinn07}. High-mass ($M > 8$ M$_{\\sun}$) and low-mass ($M <2$ M$_{\\sun}$) star formation seem to require some fundamentally different conditions and physical processes that probably involve differences in turbulence and magnetic fields within their formative clouds as much as differences in cloud mass and density \\citep{krumkee05}. Stars forming in the 2 to 8 solar mass range (intermediate-mass; IM) are especially interesting as they straddle the boundary between the low- and high-mass paradigms and can provide a window into the processes that govern the transition from low-mass to high-mass star formation. Extensive studies of these IM star-forming regions (IM SFRs), i.e., SFRs where the most massive star is IM, have not been possible in the past due to difficulties in clearly identifying the regions and the lack of angular resolution to investigate them in detail.  Our goals in this investigation of intermediate-mass star-forming regions are to understand what controls or enables the transition from low- to high-mass star formation. Is it simply a question of an enhanced reservoir of material that promotes the formation of a more massive star, or is star formation shut off in some cases because of disruption of the parent molecular cloud? Is there evidence for sequential or triggered star formation in the interstellar medium (ISM) surrounding IM SFRs as around massive SFRs \\citep{2005A&A...433..565D}? How common are bipolar outflows? What is the ISM environment of these regions? Are IM SFRs found in locations where high-mass star formation is also occurring or in relative isolation? What are the typical masses and velocity dispersions of IM star-forming clouds? What are the typical sizes and luminosities of IM SFRs?  The aim of the present paper, the first high-resolution study of IM SFRs, is to determine some basic properties of a carefully selected sample of IM SFRs. We present our candidate selection criteria and classification in Section~\\ref{section2}. We then, in Section~\\ref{section3}, identify CO features corresponding to the candidates and discuss a particularly interesting sample of the candidates. The photometry and spectral energy distribution is examined in Section~\\ref{sed}, and the molecular material is discussed in Section~\\ref{masssection}. We compare our sample to Ultra-Compact \\ion{H}{2} (UC\\ion{H}{2}) regions in Section~\\ref{uchiisection}. Finally, we present a discussion in Section~\\ref{discussion} and conclusions in Section~\\ref{conclusions}. ",
        "conclusions": "{\\label{conclusions}} We have identified a sample of 50 IM SFRs in the inner Galaxy. The IM SFRs share some basic properties. They have typical luminosities of $\\sim 10^4$ L$_{\\sun}$ and the PDRs that demarcate the IM SFRs have typical diameters of $\\sim 1$ pc. All but one lack radio continuum emission and some show small stellar clusters visible in the NIR. These properties are consistent with these objects being regions containing stars in the $2-8$ M$_{\\sun}$ range with their associated clusters of low-mass stars. The most massive stars in these regions are likely to be precursors to Herbig AeBe stars.  IM SFRs do show some similarities with high-mass star-forming regions. On parsec-scales IM SFRs are typically associated with molecular clumps of mass $\\sim 10^3$ M$_{\\sun}$ and on larger scales they are found within GMCs with masses between $10^4-10^5$ M$_{\\sun}$. We conclude that massive star formation in GMCs in that mass range is not inevitable.  IM SFRs are also distinct from regions of more massive star formation, forming in clumps with masses as low as $\\sim 10^2$ M$_{\\sun}$ and not being found in the most massive GMCs ($\\gtrsim 10^6$ M$_{\\sun}$). IM SFRs typically have an order of magnitude less luminosity than UC\\ion{H}{2} regions.  Our sample of IM SFRs, combined with UC\\ion{H}{2} region data, shows that a strong correlation between luminosity and associated molecular mass in star-forming regions applies over three orders of magnitude encompassing both IM SFRs and UC\\ion{H}{2} regions.  The molecular material associated with IM SFRs typically have an order of magnitude smaller peak mass column density and clump mass compared to UC\\ion{H}{2} regions. The lack of IM SFRs found with associated peak mass column densities $> 0.5$ g~cm$^{-2}$ supports the idea that there is a threshold in mass column density above which only massive star formation occurs."
    },
    "1005/1005.3310_arXiv.txt": {
        "abstract": "Gravitational parity violation is a possibility motivated by particle physics, string theory and loop quantum gravity. One effect of it is amplitude birefringence of gravitational waves, whereby left and right circularly-polarized waves propagate at the same speed but with different amplitude evolution. Here we propose a test of this effect through coincident observations of gravitational waves and short gamma-ray bursts from binary mergers involving neutron stars. Such gravitational waves are highly left or right circularly-polarized due to the geometry of the merger. Using localization information from the gamma-ray burst, ground-based gravitational wave detectors can measure the distance to the source with reasonable accuracy. An electromagnetic determination of the redshift from an afterglow or host galaxy yields an independent measure of this distance. Gravitational parity violation would manifest itself as a discrepancy between these two distance measurements. We exemplify such a test by considering one specific effective theory that leads to such gravitational parity-violation, Chern-Simons gravity. We show that the advanced LIGO-Virgo network and all-sky gamma-ray telescopes can be sensitive to the propagating sector of Chern-Simons gravitational parity violation to a level roughly two orders of magnitude better than current stationary constraints from the LAGEOS satellites. ",
        "introduction": "Several reasons exist to believe that there is some degree of parity violation in gravity, at least at the quantum scale and possibly carrying through to macroscopic scales through the anomaly mechanism. From an experimental standpoint, it is curious that the weak interaction exhibits maximal parity violation, yet the other fundamental forces apparently exhibit none. From a theoretical standpoint, parity violation arises from the standard model of particle physics~\\cite{Bell:1969ts,1969PhRv..177.2426A,AlvarezGaume:1983ig}, some sectors of string theory~\\cite{Green:1987mn,Alexander:2004xd}, and extensions of loop quantum gravity~\\cite{Ashtekar:1988sw,Taveras:2008yf,Calcagni:2009xz,Mercuri:2009zt,Gates:2009pt}.  Gravitational parity violation can be encoded in the action as an extension to general relativity (GR) that consists of the addition of an antisymmetric product of curvature tensors coupled to a scalar field, which to leading order in the curvature leads uniquely to Chern-Simons (CS) gravity~\\cite{Alexander:2009tp}. Jackiw and Pi~\\cite{Jackiw:2003pm} demonstrated that the CS term can be covariantly embedded with its three dimensional counterpart as a consistent modification to GR. In the standard model, there exists a radiatively generated one-loop chiral anomaly that gives rise to a CS term coupled to lepton number~\\cite{Weinberg:1996kr}. In heterotic and Type I superstring theories, the CS term is generated through the Green-Schwarz gauge anomaly-canceling mechanism~\\cite{Polchinski:1998rr}. In loop quantum gravity, the scalarization of the Barbero-Immirzi parameter coupled to fermions has been shown to lead to an effective CS action~\\cite{Ashtekar:1988sw,Taveras:2008yf,Calcagni:2009xz,Mercuri:2009zt,Gates:2009pt}. Even devoid of a specific fundamental theory, one can also show that the CS term unavoidably arises in effective field theories of inflation as one of the few non-vanishing, second-order curvature corrections to the Einstein-Hilbert action~\\cite{Weinberg:2008hq}. Another phenomenological parity-violating extension to GR has been proposed where the fundamental constants of nature (e.g.,~Newton's constant) break parity invariance~\\cite{Contaldi:2008yz}.  Most investigations of gravitational wave (GW) signatures of parity violation have focused on {\\emph{amplitude birefringence}}. That is, compared to GR, right circularly-polarized waves are enhanced or suppressed and left circularly-polarized waves are suppressed or enhanced as they propagate~\\cite{Jackiw:2003pm,Lue:1998mq,Alexander:2007kv,Yunes:2008bu,Alexander:2009tp}. Such birefringence occurs when GWs propagate on a flat or curved background. In fact, the propagation of such waves over cosmological distances has been proposed as an explanation of the baryogenesis problem in the early Universe~\\cite{Alexander:2004us} and could have observable effects on the cosmic microwave background~\\cite{Alexander:2004wk, Saito:2007kt, Contaldi:2008yz} and a stochastic GW background~\\cite{Seto:2007tn, Satoh:2007gn}. Cosmological GW amplitude birefringence has been discussed in the context of GW observations of cosmological supermassive binary black holes (BHs) with the planned Laser Interferometer Space Antenna (LISA)~\\cite{Alexander:2007kv}.  Amplitude birefringence arises in propagating modes if the parity operator does not commute with the Hamiltonian. In its most general form, the effect of the plane-wave propagator can be expressed as \\begin{equation} \\label{matrix} \\left( \\begin{array}{c} {h_{+,k}(t)} \\\\ h_{\\times,k}(t) \\end{array} \\right) = e^{-i\\omega t} \\left( \\begin{array}{rr} u & iv \\\\ -iv & u \\end{array} \\right) \\left( \\begin{array}{c} h_{+,k}(0) \\\\ h_{\\times,k}(0) \\end{array} \\right). \\end{equation} Here $\\omega$ is the GW angular frequency, $t$ is some time coordinate, $h_{+,\\times,k}$ are the Fourier components of the GW of wavenumber $k$, and $u$ and $v$ are multiplicative factors characterizing the translation invariant amplification/suppression of each Fourier mode during propagation in  time. The quantity $u$ corrects for background curvature effects: $u = 1$ in a flat-background, but it acquires redshift corrections in a cosmological background (see e.g.~\\cite{Laguna:2009re}). Given a specific modified theory, Eq.~\\eqref{matrix} can be obtained by solving the modified field equations for the propagating modes of the metric perturbation, where $v$ measures the degree of parity violation. The right- and left-circular polarizations $h_{\\RL} = (h_+ \\pm ih_\\times) / \\sqrt{2}$ then propagate as \\begin{equation} \\label{matrix2} \\left( \\begin{array}{c} h_{\\R,k}(t) \\\\ h_{\\LL,k}(t) \\end{array} \\right) = e^{-i\\omega t} \\left( \\begin{array}{rr} u + v & 0 \\\\ 0 & u-v \\end{array} \\right) \\left( \\begin{array}{c} h_{\\R,k}(0) \\\\ h_{\\LL,k}(0) \\end{array} \\right). \\end{equation} In the presence of amplitude birefringence, the eigenvalue matrix of the propagator operator remains diagonal, but with modified eigenvalues. Clearly then, right circularly-polarized waves are amplified or suppressed, while left circularly-polarized waves are suppressed or enhanced relative to GR, depending on the sign of $v$.  The amplitude birefringence discussed above and in the rest of this paper is a pure propagation effect, which accumulates over propagation distance, in contrast with  wave generation effects which do not.  As the latter cause parity violation ``at the source,'' their effect is observable at any distance~\\cite{Jackiw:2003pm,Yunes:2007ss}.  Wave generation modifications also couple to matter as well as curvature~\\cite{Alexander:2008wi}, for example leading to weakly composition dependent waveforms   in BH-NS versus NS-NS mergers. Recently,~\\cite{Sopuerta:2009iy} studied parity violation in wave generation for extreme-mass ratio inspirals, finding that indeed trajectories and the multipolar generation scheme itself encode such violation at a fundamental level. The influence of such wave generation effects on GW observables would require a full (non-extreme mass ratio) post-Newtonian and multipolar analysis, which is beyond the scope of this paper.  Instead, here we concentrate on wave propagation effects only and for simplicity adopt pure GR wave generation without modified-gravity corrections, recognizing that (i) such wave generation effects are calculable on a theory-by-theory basis and (ii) their distance-independent influence on binary phase should be easily distinguishable from any secular amplification accumulated over cosmological distances.\\optional{\\footnote{For example, depending on the theory waveform generation effects will modify the binary orbital phase and thus duration comparable to some post-Newtonian order.   On the one hand, modified waveform duration impacts overall signal fourier amplitude (through the stationary-phase fourier transform $\\tilde{h}\\propto (df/dt)^{-1/2}$).  On the other hand, gravitational wave data analysis strategies are highly sensitive to orbital phase.  Any orbital phase modification comparable  can be easily identified and extracted, to a precision comparable to the relevant post-Newtonian order. We anticipate a detailed analysis with waveform generation effects will produce a Fisher matrix that allows measurements roughly as described in the text; see Footnote \\ref{foot:Fisher:WhatWeShouldDo} for further discussion. }}  In this paper, we show that an interesting test of gravitational parity violation due to propagation effects can be made with existing instruments: the SWIFT~\\cite{Gehrels:2004am} and GLAST/Fermi~\\cite{Carson:2006af} gamma-ray satellites, and ground-based LIGO~\\cite{Abbott:2007kv} and Virgo~\\cite{Acernese:2007zze} GW detectors after ongoing upgrades. Such a test relies on the detection of coincident GW/gamma-ray burst (GRB) events, a possible progenitor of which are neutron star (NS) binary mergers, that lead to a redshift observation either from afterglow or host galaxy identification. The reasons why such events are excellent to constrain parity violation are the same as those outlined in work that proposes the use of GW/GRB events to constrain GRB parameters~\\cite{Seto:2005xb} and cosmological parameters~\\cite{Dalal:2006qt,Nissanke:2009kt,Laguna:2009re}. Due to the collimation of the jet producing the GRB, the binary's orbital angular momentum at merger must point along the line of sight to Earth, and thus the GW signal must be highly circularly-polarized. Together with the accurate sky location obtained from the GRB, this allows an accurate distance measurement from the GW signal. This distance measurement can then be compared to the purely electromagnetic distance measurement, which for typical redshifts $z \\sim 0.1$ depends only on the Hubble constant. Amplitude birefringence will manifest as a discrepancy between these two distance measurements. If no discrepancy is found, the error on the distance measurements can be used to place upper limits on the possible degree of gravitational parity violation.  The test proposed here is unique in that it employs only the radiative sector of any generic, gravitational parity violating theory. Current constraints on gravitational parity violation concentrate on a specific alternative theory (CS gravity) and explore only its stationary sector through Solar System experiments~\\cite{Alexander:2007zg, Alexander:2007vt, Smith:2007jm} and binary pulsar tests~\\cite{Yunes:2008ua}. As such, these tests constrain only specific and local deviations from gravitational parity invariance in the neighborhood of the gravitational source -- that is, near Earth in the Solar System case and near J0737-3039 in the binary pulsar case. The test we discuss here constrains {\\emph{generic parity violation}} over distances of hundreds of Mpc, along the light-cone on which the amplitude birefringent GWs propagate.  This test also compares favorably to future GW tests with LISA~\\cite{Alexander:2007kv, Yunes:2008bu}. The fractional error in a distance measurement is inversely proportional to the signal-to-noise ratio (SNR) $\\rho$. Although the latter is larger in the LISA case, the sensitivity to the parity violating effect is greater here, because this effect also depends on the GW frequency, which is much larger during NS mergers. Propagation effects should accumulate with distance propagated, and on dimensional grounds this distance should be measured in gravitational wavelengths -- that is, the effect should be proportional to $D \\, f$, where $D$ is the distance propagated and $f$ is the GW frequency. While GW/GRB events will have smaller $D$s and lower $\\rho$s than LISA supermassive BH observations (both by 1--2 orders of magnitude), $f$ is 5--6 orders of magnitude higher for the LIGO-Virgo network. Thus, the sensitivity to parity violation could be 2--3 orders of magnitude better than in the LISA case, although as we shall see later this is somewhat reduced by systematic errors.  The straw-man chosen here to compare against existing tests of gravitational parity violation is CS gravity. In this theory, the magnitude of gravitational parity violation is controlled by a length parameter (in geometric units), proportional to certain time derivatives of the CS coupling field $\\theta$. Solar system observations of the LAGEOS satellites~\\cite{Smith:2007jm} have placed the bound $\\dot{\\theta}_{0} < 2000 \\; {\\rm{km}}$ at $2\\sigma$ when $\\ddot{\\theta}_{0} = 0$. (The subscripts refer to evaluation at present, i.e.\\ zero redshift.) We have found that the above constraint can be generalized assuming $\\theta$ evolves on cosmological timescales. Allowing for both $(\\dot{\\theta}_{0}, \\ddot{\\theta}_{0})$ to be nonzero to leading order, we can generalize the Solar System constraint via the replacement rule $\\dot{\\theta}_{0} \\to |\\dot{\\theta}_{0} - \\ddot{\\theta}_{0}/H_{0}|$, which then leads to the much more powerful bound $|\\dot{\\theta}_{0} - \\ddot{\\theta}_{0}/H_{0}| < 2000 \\; {\\rm{km}}$.  Our calculations suggest that the LIGO-Virgo network will be dramatically more sensitive to parity violation, as we show in Fig.~\\ref{fig:intro} for a right-handed GW observation. The gray-shaded regions (red in the color version) show the areas of phase space that could be detected or excluded with a coincident LIGO GW/GRB observation at $2\\sigma$. This is to be compared to the regions of phase space ruled out by the generalized Solar System constraints, denoted in the figure as the region outside dotted lines. Inside the dotted lines, the generalized Solar System constraint cannot rule out the CS modification. \\begin{figure} \\includegraphics[width=\\columnwidth,clip=true]{fig-mma-ResonantIntegralSNR-Ideal-Overview} \\caption{\\label{fig:intro} Given a coincident right-handed GW/GRB detection at an SNR of $10$, the gray-shaded region (red in the color version) shows the area of $(\\dot{\\theta}_{0},\\ddot{\\theta}_{0})$ parameter space that could be excluded at $2\\sigma$ by a measurement of one right-handed short GRB merger waveform consistent with its expected GR amplitude, given host galaxy distance and (measured) chirp mass.  The region excluded by Solar System observations ($|\\dot{\\theta}_{0} - \\ddot{\\theta}_{0}/H_{0}| < 2000 \\; {\\rm{km}}$ at $2\\sigma$) is outside the dotted lines.} \\end{figure} A GW observation could place much tighter constraints on the allowed region of parity violating phase space. Such constraints are interesting because gravitational parity violation need not be Planck suppressed in the fundamental theory from which CS gravity derives. This is because the effective CS theory is shift-invariant, and thus, the scalar field that sources the correction must be massless, if this symmetry is not to be broken.  A certain asymmetry in the possible constraints we could place on $(\\dot{\\theta}_{0},\\ddot{\\theta}_{0})$ given a single, coincident GW/GRB detection is clearly discernible in Fig.~\\ref{fig:intro}. This asymmetry is an artifact of assuming a purely {\\emph{right-polarized}} GW event. In CS gravity, right-polarized signals possess a resonant behavior if $\\dot{\\theta}_{0} > 0$ for frequencies in the LIGO sensitivity band, which greatly enhances the amount of parity violation and the constraints one could place. This situation, however, reverses if instead we had assumed a left-polarized GW observation. In such a case, we would have obtained exclusion regions identical to those of Fig.~\\ref{fig:intro} but reflected under $(\\dot{\\theta}_{0},\\ddot{\\theta}_{0}) \\to (-\\dot{\\theta}_{0},-\\ddot{\\theta}_{0})$. Therefore, a two-detector observation of a right and a left-polarized GW would allow us to constrain $(\\dot{\\theta}_{0},\\ddot\\theta_{0})$ equally well for both positive and negative $\\dot{\\theta}_{0}$, thus extending the exclusion regions of Fig.~\\ref{fig:intro}.  All the results we just described assume a somewhat idealistic test, where we have ignored certain astrophysical uncertainties, such as inaccuracies in host galaxy identification. Perhaps one of the most important uncertainties relates to the inclination angle. GRB jets are not necessarily perfectly aligned with the line of sight, which implies that the inclination angle is not exactly known {\\emph{a priori}} and the GW signal is not perfectly circularly polarized. Since the inclination angle is degenerate with the luminosity distance in the GW response function, this deteriorates GW measurements of the latter by roughly a factor of two. However, even after taking several such uncertainties into account, we find that $\\dot{\\theta}_{0}$ could be constrained at the $10$-$50$~km level, which is still much better than current Solar System constraints and even LISA. Moreover, even when accounting for such uncertainties, coincident GW/GRB measurements could constrain a much larger region of the parity violating phase space all along the light-cone in which GWs propagate.  The remainder of this paper presents additional details that lead to the results described above. We begin with a description of CS gravity and amplitude birefringence in Sec.~\\ref{CS-basics} as a warm-up for the introduction of generic gravitational parity violation and its GW observable in Sec.~\\ref{gen-amp-bire}. We then proceed with a somewhat idealized test in Sec.~\\ref{Simple-test}, where we compute the modified SNR for a parity-violating signal. Assuming a perfect electromagnetic and GW determination of the luminosity distance, we then compute the constraint one could place on parity violation, assuming the latter can be treated as a linear correction to GR. In CS gravity, however, certain resonances arise in the GW signal at the non-linear level when parity violation is allowed to be strong. Such resonances greatly enhance the parity-violating corrections to the SNR, and thus, allow one to place the much stronger constraints shown in Fig.~\\ref{fig:intro}, as we show explicitly in Sec.~\\ref{res-mag}. One might worry that resonances could signal the presence of instabilities, but we show that this is not the case in the Appendix. The proposed test, however, is rather idealized because correlations with other parameters and astrophysical uncertainties generically degrade measurements of the luminosity distance. These potential sources of error are discussed in Sec.~\\ref{Ext-disc} and combined with the idealized analysis to construct realistic estimates of the possible constraints one could place on parity violation.  In the rest of this paper we employ geometric units, where the speed of light and Newton's gravitational constant are set to unity $c = 1 = G$. We further follow the conventions of Misner, Thorne and Wheeler~\\cite{Misner:1974qy}, where the metric signature is $(-,+,+,+)$, spacetime indices are denoted with Greek letters, spatial indices with Latin letters and abstract tensors are denoted in boldface type. ",
        "conclusions": "\\label{Ext-disc}  In this section we discuss estimates of various factors that improve the order of magnitude results from the previous section. We also discuss limiting noise sources that could potentially deteriorate the constraints.  \\subsection{Constraints from single short GRBs}  A network of $N_d$ GW detectors using coherent data analysis is more sensitive than one detector by a factor $\\sqrt{N_d}$ in $\\rho$, neglecting differences in orientation and noise spectrum. The four-detector LIGO-Virgo network thus gains about a factor of 2 for a typical sky position, though the full beam pattern is complicated~\\cite{Stuver:2006gr,CutlerFlanagan:1994}. Moreover, this network also is more robust than a single detector if the GRB is viewed off-axis~\\cite{Dalal:2006qt, Seto:2005xb}.  A BH/NS merger takes advantage of the mass dependence of the threshold~(\\ref{single}), $\\mathcal{M}^{5/6}$, which goes as the square root of the total mass. Thus a $30~M_\\odot$ BH ($\\mathcal{M} = 4.7~M_\\odot$) would improve the threshold further by a factor of about 4. Above $30~M_\\odot$ the improvement drops because the system starts merging before the frequency has swept through the LIGO band and because our analysis assumes no spin-orbit precession.  If the BH is rapidly spinning and that spin is misaligned with the orbit, precession will lead to a time-dependent left/right polarization ratio, reducing $\\rho$ from the on-axis case and complicating the analysis.  Spin-orbit misalignment could be large for dynamically-formed binaries. In BH-NS binaries formed from isolated binary evolution, spin and orbital angular momenta should usually be aligned and precession should be relatively unimportant; see e.g.,~Fig.~6 in \\cite{2004PhRvD..69j2002G}.  Observations and population arguments  suggest that GRB associated with NS/NS mergers have a narrow but nonzero range of opening angles; population arguments suggest narrow jets $\\sigma_{\\xi^2} = 0.02$~\\cite{2006ApJ...650..261S}, while  the lack of jet breaks in afterglows suggests that some jets may be as wide as $25^o$ \\cite{2006ApJ...653..462G} (which is still $\\xi\\approx0.9$). The finite collimation angle of the GRB will lead to a distribution of $\\xi^2$ which will degrade the threshold if $\\sigma_{\\xi^2} \\approx 1/\\rho$ since $\\sigma_{\\xi^2}$ adds roughly in quadrature with $\\sigma_D/D$.  Additionally,  because $\\xi^2 < 1$, these unknown angles can only reduce $\\rho$, not increase it, introducing a weak negative bias towards smaller amplitudes.   Without independent direct or a priori constraints on $\\xi$, any short GRB amplitude measurement by itself can only constrain the particular combination $\\lambda_{\\RL} \\left<v\\right> + (\\lambda_{\\RL} \\xi - 1)$, as \\begin{eqnarray} \\frac{\\rho}{\\hat{\\rho}(\\xi)}&\\approx& 1 + \\lambda_{\\RL} \\left<v\\right> \\left[1 - \\frac{1}{2} \\frac{F_+^2 - F_\\times^2}{F_+^2 + F_\\times^2} (\\lambda_{\\RL} - \\xi)^2 \\right. \\nonumber \\\\ &+& \\left. O(\\lambda_{\\RL}-\\xi)^3 \\right]\\,, \\end{eqnarray} where we have explicitly expanded Eq.~\\eqref{eq:amp:linear} for a nearly face on, right- or left-handed source $\\xi \\approx \\lambda_{\\RL}$ and used that $\\hat{\\rho}(\\xi) \\sim \\hat{\\rho}(\\lambda_{\\RL}) [1 + \\lambda_{\\RL} (\\xi - \\lambda_{\\RL})]$. In practice, however, the LIGO-Virgo network (or any network of differently aligned detectors) is sensitive to both GW polarizations, and by projecting directly onto a circularly-polarized basis, this network can directly constrain $1-\\xi\\simeq O(1/\\rho)$. [A similar effect was observed in~\\citet{Seto:2005xb} and~\\citet{Dalal:2006qt}, where the LIGO-Virgo network (or any network of differently aligned detectors) can reduce the systematic errors in distance measurement to short GRBs induced by wide jets.] The net effect of a narrow  distribution in $\\xi$ is therefore simply to weaken any single-observation constraint on $\\left<v\\right>$ by roughly a factor of 2, from $1/\\rho$ to $O(2)/\\rho$.  Another source of error is due to mismodeling of the signal because of the neglect of higher harmonics present for example in PN amplitude corrections~\\cite{VanDenBroeck:2006qu}. Such corrections appear in the stationary-phase of the Fourier transform as modifications to the amplitude of the form ${\\cal{A}} \\to {\\cal{A}} \\left( 1 + \\zeta_{1/2} u^{1/3} + \\zeta_{1} u^{1/3}  + \\zeta_{3/2} u \\right)$, where $u = 2 \\pi M f$ is the reduced GW frequency and $\\zeta_{n}$ are $n$th order PN amplitude coefficients, which depend on the mass ratio, the inclination angle and the polarization angles. Notice that the $1.5$PN correction, $\\zeta_{3}$, is proportional to $f$, which immediately suggests a strong-correlation and degeneracy with CS amplitude birefringence. This degeneracy should be affect the constraints quoted above much because the mass ratio is primarily determined by the phase at $1$PN order, while $<v>$ does not enter the phase at all. A full study of such possible degenerate behavior would require a full Fisher or Bayesian analysis that we shall carry out elsewhere.\\optional{\\footnote{\\label{foot:Fisher:WhatWeShouldDo}The measurement accuracy of $\\dot{\\theta}_0,\\ddot{\\theta}_{0}$ should strictly be determined using a Fisher matrix, as in \\citet{Alexander:2007kv}.  Further, this Fisher matrix should be constructed incorporating all terms comparable to and potentially correlated with the amplification discussed here, including pure GR post-Newtonian higher harmonics as well as modified-gravity wave-generation effects.  We will provide the relevant wave generation corrections and this matrix in a forthcoming paper. }}  The electromagnetic determination of the distance to the GW/GRB also has errors. The latest augmented WMAP results~\\cite{Jarosik:2010iu} imply $H_0 = 70.4\\pm1.3$~km~s$^{-1}$~Mpc$^{-1}$ at 95\\% confidence.  The intrinsic uncertainty in the Hubble constant translates to a relative distance uncertainty of  roughly   $\\delta H_{0} = 1.3/70.4\\lesssim 1.8 \\%$, far below the intrinsic SNR uncertainty $2/\\rho$ for all but  extremely bright sources ($\\rho\\gtrsim108$). Peculiar velocities of host galaxies, internal orbital velocities within galaxies and galaxy clusters, and kick velocities of binaries due to supernova kicks should all be well under $1500{\\rm km/s}\\simeq 0.005 c$ \\cite{2008NewAR..51..884M}, % corresponding to a distance error of about 6\\% in a Hubble redshift of $z=0.08$. Only for extremely bright sources or very close double NS binaries will non-cosmological velocities produce a noticeable error.  A more significant and less well controlled distance error arises through the redshift determination process itself. Currently short GRB distances can be estimated only through the redshift to electromagnetically associated hosts, and therefore depends critically on  correct host identification \\cite{GalYam:2005jf,NakarReviewArticle2006,Gehrels-shortgrb-SwiftReview-Mid2007}.  For moderate to high redshift sources, particularly where short GRBs are assumed ejected from their putative host galaxy, associations and therefore redshift measurements could be significantly confused \\cite{2009ApJ...705L.186Z}. Fortunately, for the distances to which even advanced gravitational detectors are sensitive, relatively few galaxies can conceivably be associated with a coincident gravitational and electromagnetic signature (i.e., adopting a fiducial galaxy density $n_{gal}=0.01 \\; \\unit{Mpc}^{-3}$, the number of galaxies in the electromagnetic sky-position angular extent $\\theta$ and radial extent $D$  error cone is of order $n_{gal} D^3 \\pi \\theta^2/3 = {\\cal{O}}(1)$ for $D=600\\unit{Mpc}$ and $\\theta\\simeq 1\\unit{arcmin}$).  We therefore assume any future coincident short GRB and GW signal arising within $z<0.4$ will have a correct host association and therefore a reliable redshift measurement.  Let us then summarize the enhancement and deterioration factors we have just discussed. A network of $N_{d}$ detector enhances the bound by a factor of $\\sqrt{N_{d}}$, while a BH/NS event enhances the constraint by a factor of $4$. On the other hand, an unknown jet width deteriorates the constraint by a factor of $2$, while a suboptimal search strategy also worsens the bound by a factor of $2$. Putting all of these together, we expect a modification of the constraints quoted earlier by a factor of $1/\\sqrt{N_{d}}$ for a NS/BH event or $4/\\sqrt{N_{d}}$ for a NS/NS event (i.e., from $\\left<v\\right> < 1/\\hat{\\rho}$ to $1/( \\hat{\\rho} \\sqrt{N_d})$ for a NS/BH event, where $\\hat{\\rho}$ is the single-detector SNR).  Putting all of this together, we arrive at the following constraints. A single detector with advanced LIGO sensitivity measuring a non-resonant, NS/NS inspiral signal seen directly overhead with unknown jet width leads to a $1\\sigma$ constraint of $|\\dot{\\theta}_0-\\ddot{\\theta}_0/H_0| \\leq 2400 \\; \\unit{km}$. A four-detector network would improve this limit to $|\\dot{\\theta}_0-\\ddot{\\theta}_0/H_0| \\leq 1200 \\; \\unit{km}$. A BH-NS merger, seen at comparable sensitivity over a longer baseline, would reduce the limit to $|\\dot{\\theta}_0-\\ddot{\\theta}_0/H_0| \\leq 300 \\; \\unit{km}$. All of this has assumed an optimal search strategy, but with a suboptimal one, i.e.,~using pure-GR templates, one would only reach a bound of $|\\dot{\\theta}_0-\\ddot{\\theta}_0/H_0|<600\\unit{km}$ for BH/NS events or $<2400 \\; \\unit{km}$ for NS/NS ones.  Since the maximum $(\\dot{\\theta}_0,\\ddot{\\theta}_0)$ allowed by these bounds would lead to a resonance in band, the absence of such a resonance could allow $(\\dot{\\theta}_0, \\ddot{\\theta}_0)$ to be constrained much better. In fact, this could constrain $|\\dot{\\theta}_0-\\ddot{\\theta}/H_0| \\leq 50 \\; \\unit{km}$ to $1\\sigma$ assuming an optimal search strategy (see e.g.,~Fig.~\\ref{fig:amp} for $2\\sigma$ constraints). If we additionally assume a four-detector coincident GW/GRB measurement of a NS/BH signal, the absence of a resonance could lead to a bound of\\footnote{This constraint is an optimistic one, since if this bound is saturated, the resonance occurs near $2 \\; {\\rm{kHz}}$. At such high frequencies our calculations is limited by systematic errors due to incorrect waveform modeling, since one should in principle include the plunge and merger parts of the waveform in the modeling as well.} $|\\dot{\\theta}_0-\\ddot{\\theta}_0/H_0| \\leq 25 \\; \\unit{km}$, even allowing for an unknown jet width and a suboptimal search strategy.  \\subsection{Population constraints }  Combining $N_0$ GW/GRB observations will improve the CS threshold roughly as $\\sqrt{N_0}$. Birefringence would manifest as a {\\emph{bimodal distribution}} of $\\rho/\\hat{\\rho}$, with peaks on either side of unity. If GRB jets are poorly collimated, both peaks of the distribution of $\\hat\\rho/\\rho$ would be smeared downward due to the distribution of $\\xi$ (see above).  Advanced detectors will observe very many short GRBs with GW associations. Roughly  $R_{sky}\\simeq 100-200$ short GRBs occur on the sky per year  that produce enough flux to be detected by the current generation of satellites \\cite{PSgrbs-popsyn,1999ApJS..122..465P,GuettaPiran2005}. As a first approximation, host associations suggest that the \\emph{satellite-detected} short GRB population is roughly linearly distributed in redshift out to $z\\simeq O(1)$ \\cite{GuettaPiran2006,GuettaPiran2005,PSgrbs-popsyn,NakarReviewArticle2006}.   Therefore, adopting an isotropic network horizon\\footnote{Because each short GRB is presumed optimally oriented, the range to which it can be detected is significantly greater than the angle-averaged range to which a typical source could be observed.  In practice, the distance to which short GRBs could be detected will be greater still, because when a shorter stretch of data is searched, the detection threshold can be set significantly lower~\\cite{Abbott:2005yy,Abbott:2008zzb}.}  distance $D_{gw}>200{\\rm Mpc}$, large enough to enclose many binary mergers per year,  roughly \\begin{eqnarray} R_{gw+GRB} &\\simeq& \\frac{f_{sky}}{z_{max}} D_{gw}H_0 R_{sky}, \\\\ &\\simeq&  \\frac{20}{{\\rm yr}} \\frac{f_{sky}}{z_{max}} \\frac{D_{gw}}{445 \\; {\\rm Mpc}} \\frac{R}{200 \\; {\\rm yr}^{-1}} \\nonumber \\end{eqnarray} satellite-detected  short GRBs could be detected by a GW interferometer network. Even under more conservative assumptions including the local event rate and expected beaming angles, a significant number of NS-NS or BH-NS binaries should be coincident with a short GRB event: more than  $>O(10\\%)$ of all GW detections.  Given the latest event rate predictions \\cite{O'Shaughnessy:2006wh,PSellipticals}, at least 2 should be detected per year. To summarize, a population of at least $4$ short GRBs could improve constraints on $\\dot{\\theta}_0$ by a factor $\\sqrt{4}= 2$.  Combining this factor of two and those of the previous subsection, we find that parity violation could be constrained to the level of $\\left<v\\right> < 1/(\\hat{\\rho} \\sqrt{N_d N_{0}})$ for a NS/BH event. Modeling parity violation with CS gravity and first adopting a non-resonant parity-violating event, it is reasonable to expect that after two calendar years (the length of the LIGO S5 run) an advanced LIGO-Virgo network operating in coincidence with GRB observatories would yield a 1$\\sigma$ confidence detection of or upper limit on $|\\dot{\\theta}_0-\\ddot{\\theta}_o/H_0|$ of $150 \\; \\unit{km}$ for BH-NS ($600 \\; \\unit{km}$ for NS-NS) using optimally matched templates, or $300 \\; \\unit{km}$ for BH-NS ($1200 \\;  \\unit{km}$ for NS-NS) using a conventional pure-GR data analysis strategy. Again, as the maximum $(\\dot{\\theta}_0,\\ddot{\\theta}_0)$ allowed by the aforementioned constraints would lead to a resonance in band, the absence of resonances could allow a much better constraint. In the neighborhood of $\\ddot{\\theta}_{0} = 0$, the absence of such resonance could yield a bound on $\\dot{\\theta}_{0}$ of ${\\cal{O}}(10 \\; \\unit{km})$.  The projected constraints on CS gravity can be thus more than an order of magnitude better than current Solar System bounds placed with the LAGEOS satellites~\\cite{Smith:2007jm}. Although none of these constraints are competitive with the one recently placed with the double binary pulsar~\\cite{Yunes:2008ua}, the latter only samples parity violation in the neighborhood of this binary, while the Solar System one is sensitive to parity violation only in the vicinity of the Earth-Moon system. The test proposed here would allow for constraints of the dynamical sector of parity violating theories along the GW geodesic, thus exploring a much larger region of the Universe.  \\subsection{Temporal-Spatial Variability of Parity Violation}  In this paper, we have proposed limits on $\\dot\\theta_0$ and $\\ddot\\theta_0$ assuming (i) that $\\dot{\\theta}$ varies on cosmological timescales and (ii) homogeneously and isotropically throughout the Universe. In principle, neither assumption need apply, as the $\\theta$ field that parametrizes parity violation in CS gravity should be prescribed by the more fundamental theory of which CS gravity is only an effective model. If the $\\theta$ field does not obey the symmetries of the background (homogeneity and isotropy in the cosmological case), then the background metrics employed, on which GW propagate, need not be solutions to the effective theory. More general $\\theta$ fields would possibly require a re-analysis of the appropriate cosmological solutions and the propagation of GWs in them.  With this caveat in mind, the more generic expressions of Eq. (\\ref{generalv}) that parameterize generic parity violating effects via the amplification process of Sec.~\\ref{res-mag} suggest that the proposed constraints are primarily on a ratio of $\\calF_{\\RL}$, with one factor at the source (the numerator) and another, more critical one at the detector (the denominator).  These two factors can be separated.  For example,  the resonant frequency and resonant bandwidth in our detectors are set primarily by the present-day evolution of $\\theta$.    If a resonance is indeed in band, \\emph{every} source with some amount of resonant polarization will be amplified roughly the same way; stacking power from multiple detections could constrain the presence of a universal spectral feature and potentially directly constrain $\\dot{\\theta}_0$, independent of its cosmological behavior, to lie outside the wavelength region to which LIGO is sensitive. On the other hand, the difference between the values of $\\dot{\\theta}$ at present and the source determine how close a zero in the numerator is to the resonance and  thus the amplification ${\\calAmp}_{\\RL}$; cf. Eq. (\\ref{eq:Amp:Approx}).  Therefore, if the detected SNRs do exhibit some  amplifications $\\rho/\\hat{\\rho}$ versus $D$ that are \\emph{not} roughly uniformly increasing with distance, these measurements can be fit to either a functional form (if the points trace a curve) or a distribution (if the points have excess scatter) for $\\dot{\\theta}$.  As a concrete example, so long as resonances are at frequencies far higher than in band ($\\kappa\\theta'\\ll 1$) and the source is at low enough redshift that cosmological terms can be ignored ($a=a_s\\approx 1$), the relative amplification  in SNR due to CS birefringence is [cf. Eq. (\\ref{eq:v:Simple2})] \\begin{eqnarray} \\rho/\\hat{\\rho}&\\approx& 1  + \\lamRL \\pi \\left<f\\right> (\\theta'_0 -\\theta'_s), \\end{eqnarray} which applies to any $\\theta'(\\eta)$ which varies on scales faster than cosmological timescales. This would then allow for generic tests of gravitational parity violation, irrespective of the specific theoretical model.   \\optional{\\subsection{Tests without EM distances or circular polarization} The test described above is straightforward precisely because it relies on special input: electromagnetically-determined distances to sources assumed to emit circularly polarized gravitational waves. More sophisticated data analysis strategies can relax both conditions.  For example, even without a known distance, the population of nearly circularly polarized sources with known sky location naturally yields a bimodal distribution of SNR, with a bias proportional to a (distance-independent) $\\delta \\rho$; the distribution of left and right handed sources therefore always constrains $\\delta \\rho$ and therefore birefringence.  Likewise, all binary sources have a mixture of left- and right-handed signal power; birefringence implies a network will recover more of one handedness than another, overall.  The relative proportion of left- versus right-handed power therefore provides another test.}"
    },
    "1005/1005.1123_arXiv.txt": {
        "abstract": "We investigate seed magnetic field generation in the early universe by radiation force of first stars. In a previous study with the steady assumption, large amplitudes($\\sim 10^{-15}$G for first stars, $\\sim 10^{-11}$G for QSOs) are predicted. In this study, we formulate this issue in an unsteady framework. Then, we consider a specific model of magnetic field generation around a very massive first star. Consequently, we (1) find that the steady assumption is not valid in realistic situations, and (2) obtain much smaller magnetic field strength than predicted by Langer et al. In addition, we find that the momentum transfer process during photoionization is more important than Thomson scattering. The resultant magnetic flux density around the first star is $\\la 10^{-19}$G. This seed magnetic field will not affect subsequent star formation in the neighborhood of first stars. ",
        "introduction": "One of the most important current issues in cosmology is understanding the star formation history of the universe, especially the early universe. Owing to progress in the last decade, mainly due to theoretical studies, we now expect the very first stars in the early universe to be very massive ($\\ga 100M_\\odot$), compared to those in present-day galaxies\\citep[e.g.,][]{Bromm02,NU01,Abel02, Yoshida06a}.  The chief reason that the first stars are very massive is the lack of heavy elements in the primordial gas, since the heavy elements are efficient coolants. Theoretical studies suggest that when the metallicity of the star forming gas exceeds $10^{-5} {\\rm to} 10^{-6}Z_\\odot$, the star formation mode transits from very massive star formation to normal star formation\\citep{schneider03,omukai05}. Thus, metallicity is the key factor for the transition of the star formation mode.  Another possible parameter is magnetic field strength. In the present-day star-forming molecular cloud, energy density of the magnetic field is comparable to gravitational or turbulent energy density, while it is assumed that the magnetic field is weak in the primordial gas. Magnetic field in present-day galaxies basically suppresses the star formation process.  The effects of magnetic field on the early star formation episode in the universe have been studied by several authors. The coupling between magnetic field and pure primordial gas has been investigated\\citep{maki04,maki07}. They found that the magnetic field is always frozen in the gas in so far as the magnetic field strength is not strong enough to prevent gravitational contraction. Based on the frozen-in assumption, it is found that bipolar outflows emerge in cases $B > 10^{-9}$G is satisfied at $n_{\\rm H}=10^3 {\\rm cm^{-3}}$ \\citep{machida06}. \\citet{tan_blackman} suggest that magnetorotational Instability(MRI) is activated in the accretion disk surrounding proto first stars in cases where $B > 10^{-10}$G at $n_{\\rm H}=10^3 {\\rm cm^{-3}}$, which results in efficient angular momentum transfer in the disk. It is also reported by \\citet{schleicher} that gravitational contraction of the primordial gas is delayed in cases where the comoving field strength is larger than $10^{-10}$G.  Therefore, it is crucial to determine the magnetic field strength that exists in the environment of the star-forming gas cloud in the early universe in order to understand the transition of the star formation mode. Magnetic field strength at the star formation site in the early universe is basically provided by the field generated prior to nonlinear structure formation \\citep[e.g.,][]{harrison1,MR72,turner,ichiki}, as well as the field generated by the Biermann battery effect\\citep{biermann} during the structure formation \\citep[e.g.,][]{xu} or reionization \\citep{gnedin}. The predicted field strength in these models is within the  range of $10^{-16}{\\rm G} {\\rm to} 10^{-20}{\\rm G}$, which is too weak to affect star formation. Another possibility is that the magnetic field spewed out from the formed stars \\citep[e.g.,][]{BRS73}. In cases where the field strength is as large as $\\sim$1 G like normal stars, a simple flux freezing assumption again leads to $\\sim 10^{-16}$G at intergalactic medium(IGM) densities.  On the other hand, \\citet{langer} suggest a mechanism to generate a large amplitude of seed magnetic field as large as $10^{-11}$G, which could affect the star formation in the early universe. They consider HI\\hspace{-.1em}I regions formed by luminous sources like QSOs or first stars. The radiation force field due to the photons emitted by these objects is anisotropic because of the clumpiness of surrounding media. This anisotropy causes charge separation in the HI\\hspace{-.1em}I region, which generates a large scale eddy current as well as the magnetic fields. However, the assumption in this theory is idealized. In their theory, anisotropy of radiation is dealt with as perturbation, and the HI\\hspace{-.1em}I region is assumed to be steady. In reality, dense shadows are generated behind the nearby dense halos. In addition, HI\\hspace{-.1em}I regions formed by QSOs or first stars are actually unsteady, and ionization fronts never reach the Str$\\ddot{\\rm o}$mgren radius \\citep{shapiro,kitayama,yoshida07}.  In this paper, we investigate this magnetic field generation process by radiation force in nonlinear and unsteady framework along the lines of theory of \\citet{langer}. We consider a specific model of magnetic field generation around a very massive first star and assess whether the generated magnetic field affects the subsequent star formation.  In Section 2, we describe the basic equations. A brief description and the setup of our models are described in Section 3. In Section 4, we show the results of our calculation, and Section 5 is devoted to discussion. Finally, in Section 6, we conclude our research. ",
        "conclusions": "We investigate the magnetic field generation process by radiation force in a nonlinear and unsteady framework along the line of the theory of \\citet{langer}. We consider a specific model of magnetic field generation around a very massive first star. Using analytical and numerical calculations, we find that the steady assumption is not valid in realistic situations and obtain a much smaller magnetic field strength than predicted by the steady theory. In addition, we find that the momentum transfer process during photoionization is more important than Thomson scattering. The resultant magnetic flux density around first stars is $\\sim 10^{-19}$G. This seed magnetic field does not seem to affect subsequent star formation in the neighborhood of first stars.  \\bigskip We thank the anonymous referee for critical comments which helped to improve the manuscript. We also thank N.Tominaga for fruitful discussions. This work was supported in part by the Inamori Foundation, as well as Ministry of Education, Science, Sports and Culture, Grant-in-Aid for Scientific Research (C), 22540295."
    },
    "1005/1005.1912_arXiv.txt": {
        "abstract": "We derive fundamental, structural, and photometric parameters of 15 overlooked Ruprecht (hereafter Ru) star clusters by means of 2MASS photometry and field-star decontamination. Ru\\,1, 10, 23, 26, 27, 34, 35, 37, 41, 54, 60, 63, 66, and 152 are located in the third Galactic quadrant, while Ru\\,174 is in the first. With the constraints imposed by the field-decontaminated colour-magnitude diagrams (CMDs) and stellar radial density profiles (RDPs), we derive ages in the range 400\\,Myr --- 1\\,Gyr, except for the older Ru\\,37, with $\\sim3$\\,Gyr. Distances from the Sun are within $\\rm1.5\\la\\ds(kpc)\\la8.0$. The RDPs are well-defined and can be described by a King-like profile for most of the radial range, except for Ru\\,23, 27, 41, 63, and 174, which present a conspicuous stellar density excess in the central region. The clusters dwell between (or close to) the Perseus and Sagittarius-Carina arms. We derive evidence in favour of cluster size increasing with distance to the Galactic plane ($\\zgc$), which is consistent with a low frequency of tidal stress associated with high-$|\\zgc|$ regions. The clusters are rather faint even in the near-infrared, with apparent integrated \\jj\\ magnitudes within $6.4\\la m_J\\la9.8$, while their absolute magnitudes are $-6.6\\la M_J\\la-2.6$. Extrapolation of the relation between $M_V$ and $M_J$, derived for globular clusters, suggests that they are low-luminosity optical clusters, with $-5\\la M_V\\la-1$. ",
        "introduction": "\\label{Intro}  In general, open clusters (OCs) are formed and evolve in - or close to - the Galactic disk. As a consequence of their orbits, they are constantly suffering tidal stress from Galactic substructures and undergoing different degrees of mass loss that, in most cases, might lead to cluster dissolution. Indeed, the vast majority of the OCs do not survive even the embedded phase (e.g. \\citealt{LL2003}), and few reach ages older than $\\sim10^8$\\,yr (e.g. \\citealt{GoBa06}).  As systems in which the balance between velocity dispersion and escape velocity plays a vital role, OC stellar distributions change continually as a function of time because of mass segregation and evaporation, tidal interactions with the disk and/or bulge, encounters with giant molecular clouds, as well as mass loss associated with stellar evolution. On average, these processes tend to accelerate the cluster dynamical evolution and change the internal structure in varying degrees. Indeed, near the solar circle, theoretical and observational evidence (e.g. \\citealt{Spitzer58}; \\citealt{Oort58}; \\citealt{BM03}; \\citealt{GoBa06}; \\citealt{LG06}; \\citealt{Khalisi07}; \\citealt{Piskunov07}) suggest a mass dependent disruption time scale of a few $10^8$\\,yr. As a consequence, most OCs end up completely dissolved in the Galactic stellar field (e.g. \\citealt{Lamers05}) or leave only poorly-populated remnants (\\citealt{PB07} and references therein), long before reaching $\\sim1$\\,Gyr of age.  Reflecting the above age/dissolution effect, only $\\approx13\\%$ of the $\\approx1100$ OCs with known age listed in WEBDA\\footnote{\\em obswww.univie.ac.at/webda} are older than 1\\,Gyr, while $\\approx45\\%$ have an age between 100\\,Myr and 1\\,Gyr, and $\\approx42\\%$ are younger than 100\\,Myr. So, besides the obvious importance of deriving reliable astrophysical parameters for as yet unstudied clusters, the unambiguous characterisation of OCs older than several $10^8$\\,yr will increase the statistics of objects undergoing the dissolution phase. This, in turn, can be used for better determining the time scale for cluster dissolution in the Galaxy.  Cluster databases such as WEBDA still contain many unstudied objects that, over the years, have been identified as star cluster candidates, usually based on the appearance in optical images. Not surprisingly, close investigations of some of these overlooked objects have uncovered a number of actual OCs (e.g. \\citealt{Carr05}; \\citealt{Carr06a}; \\citealt{Carr06b}). Besides, since young OCs are relatively easy to identify (because of the presence of bright stars) even at relatively large distances, the overlooked clusters, in general, are as old as several $10^8$\\,yr, sometimes reaching a few Gyr (e.g. the $\\sim4$\\,Gyr old OC Berkeley\\,56 - \\citealt{Carr06b}).  \\begin{figure*} \\begin{minipage}[b]{0.33\\linewidth} \\includegraphics[width=\\textwidth]{Ru1_5x5.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.33\\linewidth} \\includegraphics[width=\\textwidth]{Ru10_5x5.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.33\\linewidth} \\includegraphics[width=\\textwidth]{Ru23_5x5.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.33\\linewidth} \\includegraphics[width=\\textwidth]{Ru26_5x5.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.33\\linewidth} \\includegraphics[width=\\textwidth]{Ru27_5x5.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.33\\linewidth} \\includegraphics[width=\\textwidth]{Ru34_5x5.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.33\\linewidth} \\includegraphics[width=\\textwidth]{Ru35_5x5.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.33\\linewidth} \\includegraphics[width=\\textwidth]{Ru37_5x5.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.33\\linewidth} \\includegraphics[width=\\textwidth]{Ru41_5x5.eps} \\end{minipage}\\hfill \\caption[]{All panels display $5\\arcmin\\times5\\arcmin$ DSS-I B images. From left to right: Ru\\,1, 10, and 23 (top); Ru\\,26, 27, and 34 (middle); Ru\\,35, 37 (with the bright K\\,3\\,III star 6\\,Pup at $\\approx2\\arcmin$ to the northwest), and Ru\\,41 (bottom). Plate flaws show up in the image of Ru\\,41. Orientation: North at the top and east at left.} \\label{fig1} \\end{figure*}  \\begin{figure*} \\begin{minipage}[b]{0.33\\linewidth} \\includegraphics[width=\\textwidth]{Ru54_5x5.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.33\\linewidth} \\includegraphics[width=\\textwidth]{Ru60_5x5.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.33\\linewidth} \\includegraphics[width=\\textwidth]{Ru63_5x5.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.33\\linewidth} \\includegraphics[width=\\textwidth]{Ru66_5x5.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.33\\linewidth} \\includegraphics[width=\\textwidth]{Ru152_5x5.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.33\\linewidth} \\includegraphics[width=\\textwidth]{Ru174_5x5.eps} \\end{minipage}\\hfill \\caption[]{Same as Fig.~\\ref{fig1} for the remaining clusters. From left to right: Ru\\,54, 60, and 63 (top); Ru\\,66, 152, and 174 (bottom). Plate flaws show up in the image of Ru\\,54.} \\label{fig2} \\end{figure*}  Based on the optical {\\em Catalogue of Star Clusters and Associations} of \\citet{Alter70}, WEBDA lists the coordinates of 171 OC candidates originally found by J. Ruprecht (classified as Ruprecht clusters; hereafter Ru). However, only 79 have been so far confirmed as OCs, having determinations of fundamental parameters, such as the age, reddening, and distance from the Sun. In the present paper we derive fundamental and structural parameters for 15 of these overlooked Ruprecht OCs, 12 of these not previously studied, and the remaining 3 with inconsistent determinations (Sects.~\\ref{DFP} and \\ref{struc}). Steps taken in the present work can be summarised as follows: {\\em (i)} 2MASS\\footnote{The Two Micron All Sky Survey, All Sky data release (\\citealt{2mass1997}) - {\\em http://www.ipac.caltech.edu/2mass/releases/allsky/}} photometry is extracted in a wide circular region centred on a given cluster, {\\em (ii)} we apply field-star decontamination to uncover the intrinsic colour-magnitude diagram (CMD) morphology, which is essential for a proper derivation of reddening, age, and distance from the Sun, and {\\em (iii)} we apply colour-magnitude filters to produce more contrasted  stellar radial density profiles (RDPs). In particular, field-star decontaminated CMDs constrain the fundamental parameters more than the raw (observed) photometry, especially for low-latitude and/or bulge-projected OCs (e.g. \\citealt{ProbFSR}, and references therein).   \\begin{table*} \\caption[]{Fundamental parameters} \\label{tab1} \\tiny \\renewcommand{\\tabcolsep}{1.1mm} \\renewcommand{\\arraystretch}{1.25} \\begin{tabular}{cccccccccccccccc} \\hline\\hline &\\multicolumn{6}{c}{WEBDA}&&\\multicolumn{8}{c}{This work}\\\\ \\cline{2-7}\\cline{9-16} Cluster&$\\alpha(2000)$&$\\delta(2000)$&Age&\\ebv&\\ds&D&&$\\alpha(2000)$&$\\delta(2000)$& $\\ell$&$b$&Age&\\ebv&\\ds&\\dSC\\\\ & (hms)&($\\degr\\,\\arcmin\\,\\arcsec$)&(Myr)&(mag)&(kpc)&(\\arcmin)&&(hms)&($\\degr\\,\\arcmin\\,\\arcsec$) &(\\degr)&(\\degr)&(Myr)&(mag)&(kpc)&(kpc)\\\\ (1)&(2)&(3)&(4)&(5)&(6)&(7)&&(8)&(9)&(10)&(11)&(12)&(13)&(14)&(15)\\\\ \\hline Ru\\,1  &06:36:25&$-$14:10:48&580&0.15&1.1&5.0&&06:36:21.63&$-$14:08:48.75&223.95&$-$9.68&$500\\pm100$&$0.26\\pm0.06$&$1.72\\pm0.16$&$1.31\\pm0.12$\\\\ Ru\\,10 &07:06:25&$-$20:05:00&---&---&--- &4.0&&07:06:29.05&$-$20:06:31.50&232.55&$-$5.85&$500\\pm100$&$0.64\\pm0.06$&$2.34\\pm0.22$&$1.62\\pm0.14$ \\\\ Ru\\,23 &07:30:41&$-$23:23:00&---&---&--- &4.0&&07:30:38.80&$-$23:23:36.00&238.08&$-$2.40&$600\\pm100$&$0.54\\pm0.06$&$3.06\\pm0.29$&$2.00\\pm0.16$ \\\\ Ru\\,26 &07:37:11&$-$15:38:59&30&0.10&1.4&24.0&&07:37:11.10&$-$15:39:46.50&232.06&$+$2.68&$400\\pm50$ &$0.35\\pm0.06$&$1.82\\pm0.17$&$1.24\\pm0.11$ \\\\ Ru\\,27 &07:37:30&$-$26:31:42&250&0.15&0.6&21.6&&07:37:40.89&$-$26:31:45.50&241.60&$-$2.53&$900\\pm100$&$0.03\\pm0.06$&$1.49\\pm0.14$&$0.82\\pm0.07$ \\\\ Ru\\,34 &07:45:55&$-$20:23:00&---&---&--- &6.0&&07:45:56.22&$-$20:23:24.00&237.20&$+$2.16&$1000\\pm100$&$0.00\\pm0.06$&$2.63\\pm0.25$&$1.70\\pm0.14$ \\\\ Ru\\,35 &07:46:13&$-$31:17:00&---&---&--- &1.5&&07:46:13.64&$-$31:16:47.25&246.66&$-$3.25&$400\\pm100$&$0.45\\pm0.10$&$3.91\\pm0.56$&$2.26\\pm0.28$ \\\\ Ru\\,37 &07:49:54&$-$17:17:00&---&---&--- &4.0&&07:49:47.53&$-$17:14:46.50&234.94&$+$4.53&$3000\\pm1000$&$0.00\\pm0.06$&$5.25\\pm0.74$&$3.87\\pm0.45$ \\\\ Ru\\,41 &07:53:49&$-$26:58:00&---&---&--- &2.0&&07:53:48.38&$-$26:57:39.00&243.78&$+$0.37&$700\\pm100$&$0.13\\pm0.10$&$3.15\\pm0.45$&$1.84\\pm0.23$ \\\\ Ru\\,54 &08:11:21&$-$31:57:00&---&---&--- &3.0&&08:11:20.92&$-$31:56:49.50&250.03&$+$0.96&$800\\pm100$&$0.13\\pm0.10$&$5.47\\pm0.78$&$3.22\\pm0.43$ \\\\ Ru\\,60 &08:24:27&$-$47:13:00&---&---&--- &3.0&&08:24:26.29&$-$47:12:51.01&264.10&$-$5.51&$400\\pm100$&$0.64\\pm0.10$&$6.16\\pm0.88$&$2.74\\pm0.54$ \\\\ Ru\\,63 &08:32:40&$-$48:18:00&---&---&--- &3.0&&08:32:39.60&$-$48:18:19.50&265.80&$-$5.02&$500\\pm100$&$0.61\\pm0.10$&$3.76\\pm0.37$&$1.16\\pm0.17$ \\\\ Ru\\,66 &08:40:33&$-$38:04:00&---&---&--- &2.0&&08:40:33.82&$-$38:04:47.99&258.49&$+$2.28&$600\\pm100$&$0.90\\pm0.10$&$3.76\\pm0.54$&$1.56\\pm0.24$ \\\\ Ru\\,152&07:54:30&$-$38:14:00&---&---&--- &3.0&&07:54:28.35&$-$38:14:14.26&253.54&$-$5.30&$600\\pm100$&$0.67\\pm0.10$&$8.02\\pm1.15$&$5.00\\pm0.73$ \\\\ Ru\\,174&20:43:30&$+$37:03:00&---&---&--- &8.0&&20:43:30.48&$+$37:01:26.26& 78.02&$-$3.39&$800\\pm100$&$0.32\\pm0.06$&$2.11\\pm0.20$&$-0.13\\pm0.07$ \\\\ \\hline \\end{tabular} \\begin{list}{Table Notes.} \\item Col.~7: diameter estimated from optical images (\\citealt{diasCat}); Col.~14: distance from the Sun; col.~15: distance from the Solar circle. \\end{list} \\end{table*}  This paper is organised as follows. In Sect.~\\ref{RecAdd} we present optical images of the sample objects. In Sect.~\\ref{2mass} we discuss the 2MASS photometry and build the field-star decontaminated CMDs. In Sect.~\\ref{DFP} we derive fundamental cluster parameters. In Sect.~\\ref{struc} we derive structural parameters. In Sect.~\\ref{Discus} we investigate relations among parameters and with respect to their location in the Galaxy. Concluding remarks are given in Sect.~\\ref{Conclu}. ",
        "conclusions": "\\label{Conclu}  Given the rather efficient and numerous dissolution mechanisms operating in the Galaxy, the majority of the open clusters do not survive beyond a few $10^8$\\,yr. In this context, it is important to investigate the structural and photometric properties of OCs that are undergoing this evolved phase.  The present paper focuses on 15 overlooked Ruprecht clusters, 12 of them never before studied. With the exception of a single object at $\\ell\\approx78\\degr$, the remaining clusters are located in the $3^{rd}$ Galactic quadrant, which minimises the field-star contamination. We work with 2MASS photometry (with errors $\\la0.15$\\,mag), on which we apply field-star decontamination to enhance CMD evolutionary sequences and stellar RDPs, thus yielding more constrained fundamental and structural parameters.  As typical optically discovered clusters, the reddening values are relatively low, within $0.0\\la\\ebv\\la0.9$; on the other hand, they are distant from the Sun, within $\\rm1.5\\la\\ds(kpc)\\la8.0$. The integrated apparent \\jj\\ magnitudes are rather faint, within $6.4\\la m_J\\la9.8$, but given the distances, the absolute magnitudes are relatively bright $-6.6\\la M_J\\la-2.6$. The ages are in the range 400\\,Myr - 1\\,Gyr, except for the significantly older Ru\\,37, with $\\sim3$\\,Gyr. The RDPs are well contrasted with respect to the background and follow the King-like profile for most of the radial range. Exceptions are Ru\\,23, Ru\\,27, Ru\\,41, Ru\\,63, and Ru\\,174, which present a pronounced stellar density excess in the innermost RDP bin. The core radii of the present sample are small, when compared to those of nearby OCs (e.g. \\citealt{Piskunov07}). By extrapolating the relation between $M_V$ and $M_J$, derived for globular clusters, we estimate $-5\\la M_V\\la-1$, which suggests that they are low-luminosity optically-selected clusters.  The sample clusters are located between (or close to) the Perseus and Sagittarius-Carina arms, and we present evidence that the cluster size increases both with Galactocentric distance and distance to the plane. The latter relation, in particular, is consistent with a low frequency of tidal stress associated with high-$|\\zgc|$ regions.  It is clear from the above analysis that searches for star clusters in catalogues of candidates - even in the optical - are far from complete. Detailed investigations will certainly add more members to the present-day open cluster census. Besides, since young clusters are rather easy to identify even at large distances, the overlooked clusters are expected to be of the evolved/old age range. Thus, works like the present one are important not only because reliable astrophysical parameters are derived for a sample of unstudied clusters. Perhaps the main importance lies in the unambiguous characterisation of open clusters with ages beyond several $10^8$\\,yr. A better statistics on the population of these - and older - clusters can be used to investigate cluster formation rates and to constrain the time scale of cluster dissolution in the Galaxy."
    },
    "1005/1005.0971_arXiv.txt": {
        "abstract": "We present the first results of our program to study a sample of local luminous infrared galaxies (LIRGs, $L_{\\rm IR}$ = 10$^{11}$-10$^{12}~L_{\\rm \\odot}$) with the \\textit{Spitzer} infrared spectrograph (IRS). In these proceedings we investigate the behavior of the 9.7 \\micron\\ silicate feature in LIRGs. As opposed to the extreme silicate absorptions observed in ultraluminous infrared galaxies (ULIRGs, $L_{\\rm IR}$ = 10$^{12}$-10$^{13}~L_{\\rm \\odot}$), LIRGs exhibit intermediate silicate absorption features, comparable to those of starburst galaxies. We also find that most of the LIRGs have the minima of the [NeIII]\\slash[NeII] ratio located at their nuclei. It is likely that increased densities in the nuclei are responsible for the smaller nuclear ratios. In the nuclei, it is also possible that the most massive stars are either absent, or still embedded in ultracompact HII regions. Finally we discuss the possible contribution of an AGN to the nuclear mid-IR emission of the galaxy, which in general is low in these local LIRGs. ",
        "introduction": "Luminous infrared galaxies (LIRGs, $L_{\\rm IR}$ = 10$^{11}$-10$^{12 } L_{\\rm \\odot}$) are an important cosmological class of galaxies as they are the main contributors to the co-moving star formation rate density of the universe at $z$ = 1. Moreover, the mid-IR spectra of high redshift (z $\\sim$ 2) very luminous IR galaxies ($L_{\\rm IR}$ $>$ 10$^{12} L_{\\rm \\odot}$) appear to be better reproduced by those of local starbursts and LIRGs.  We obtained \\textit{Spitzer}/IRS mapping observations and analyzed archival IRS staring observations of the sample of local (d $<$ 75 Mpc) LIRGs of \\citet{AAH06s_mps}. The maps cover the central 20$\\arcsec\\times$20$\\arcsec$ or 30$\\arcsec\\times$30$\\arcsec$ regions of the galaxies, and use all four IRS modules to cover the full 5-40 $\\micron$ spectral range. The IRS spectral mapping data reduction is described in detail by \\citet{MPS2010_mps}. Briefly, we assembled the spectral data cubes from the staring observations with CUBISM. Then we used the low spectral resolution (R $\\sim$ 60-120) data cubes to measure the PAH and the 9.7 \\micron\\ silicate features and we obtained maps of the brightest emission lines by fitting the high resolution (R $\\sim$ 600) spectra. Using these emission line maps we calculated ratio maps. ",
        "conclusions": ""
    },
    "1005/1005.2974_arXiv.txt": {
        "abstract": "The aim of this study is to search for observational evidence of vertical iron stratification in the atmosphere of fourteen blue horizontal-branch (BHB) stars. We have found from our numerical simulations that five BHB stars: B22, B186 in the globular cluster NGC~288, WF2-820, WF2-2692 in M13 and B203 in M15 show clear signatures of the vertical stratification of iron whose abundance increases toward the lower atmosphere. Two other BHB stars (B334 in M15 and B176 in M92) also show possible iron stratification in their atmosphere. A dependence of the slope of iron stratification on the effective temperature was also discovered. It is found that the vertical stratification of iron is strongest in BHB stars with $T_{\\rm eff}$ around 11,500K. The slope of iron abundance decreases as $T_{\\rm eff}$ increases and becomes negligible for the BHB stars with $T_{\\rm eff} \\approx$ 14,000K. These results support the hypothesis regarding the efficiency of atomic diffusion in the stellar atmospheres of BHB stars with $T_{\\rm eff}\\succeq$ 11,500~K. ",
        "introduction": "According to the current understanding of stellar evolution, the horizontal-branch (HB) stars are post-main sequence stars that burn helium in their core and hydrogen in a shell (e.g. Hoyle\\& Schwarzschild 1955; Moehler 2004). In this paper we consider the HB stars that are located in the blue part of the HB, to the left of the RR Lyrae instability strip. These stars are commonly called blue horizontal-branch (hereafter BHB) stars to distinguish them from the red horizontal-branch (RHB) stars, which exhibit different observational properties.  \\begin{table*}% \\centering \\begin{minipage}{150mm} \\caption[]{Journal of Keck+HIRES spectroscopic observations of the selected hot BHB stars from Behr \\shortcite{Behr03a}.} \\label{tab1} \\begin{tabular}{@{}lccccccccccc@{}} \\hline Cluster/Star&     HJD    &Exposure&  S/N & Coverage &$T_{\\rm eff}$&$\\log{g}$&$V\\sin{i}$ &$V_{\\rm r}$& [Fe/H] & [He/H]\\\\ &2450000+&   Time (s)   & &  (\\AA)  &   (K)     &  (dex)    &(km s$^{-1}$)&(km s$^{-1}$)& & \\\\ \\hline M13/WF2-820 & 1052.7815 &3$\\times$1200& 34 & 3888-5356 & 11840 & 3.90 & 4.17 &-239.7$\\pm$0.7&-0.22$^{+0.14}_{-0.23}$&-1.09$^{+0.47}_{-0.44}$\\\\ M13/WF2-2692& 1047.7400 &3$\\times$1800& 34 & 3885-6292 & 12530 & 4.08 & 4.64 &-236.2$\\pm$0.7&-0.23$^{+0.76}_{-0.34}$&-1.81$^{+0.96}_{-1.07}$\\\\ M13/WF2-3123& 1248.1298 &3$\\times$1200& 35 & 3885-6292 & 13670 & 4.31 & 6.32 &-238.3$\\pm$1.1&-0.71$^{+0.65}_{-0.51}$&-1.28$^{+1.55}_{-0.92}$\\\\ M13/WF3-548 & 1053.7293 &3$\\times$1200& 32 & 3888-5356 & 13100 & 4.16 & 4.29 &-235.3$\\pm$0.7&-0.64$^{+0.60}_{-0.35}$&-2.10$^{+1.08}_{-0.98}$\\\\ M13/WF3-1718& 1248.0811 &3$\\times$1200& 48 & 3888-5884 & 11510 & 3.79 & 1.89 &-244.0$\\pm$0.3&+0.02$^{+0.10}_{-0.13}$&-1.86$^{+0.42}_{-0.68}$\\\\ M15/B203    & 1408.0180 &3$\\times$1200& 28 & 3885-6292 & 13990 & 3.84 & 4.89 & -94.7$\\pm$1.2&+0.02$^{+0.38}_{-0.42}$&-2.54$^{+1.15}_{-0.67}$\\\\ M15/B315    & 1409.0238 &3$\\times$1200& 24 & 3885-6292 & 12890 & 3.81 & 1.72 &-103.3$\\pm$0.8&+0.04$^{+0.44}_{-0.36}$&-1.75$^{+0.94}_{-0.99}$\\\\ M15/B334    & 1052.8791 &4$\\times$1500& 39 & 3888-5356 & 10750 & 3.61 & 9.22 &-109.3$\\pm$1.3&-2.37$^{+0.19}_{-0.21}$&-0.64$^{+0.45}_{-0.45}$\\\\ M15/B374    & 1406.0183 &3$\\times$1200& 21 & 3885-6292 & 12820 & 3.82 & 3.92 &-106.5$\\pm$0.9&+0.24$^{+0.48}_{-0.37}$&-1.97$^{+0.94}_{-0.92}$\\\\ M92/B176    & 1049.0646 &3$\\times$1800& 56 & 4020-5520 & 11150 & 3.76 & 6.96 &-116.3$\\pm$1.4&-2.25$^{+0.16}_{-0.24}$&-0.55$^{+0.36}_{-0.32}$\\\\ NGC288/B16  & 1406.0850 &3$\\times$1200& 21 & 3885-6292 & 14030 & 4.15 & 3.08 & -44.7$\\pm$0.9&+0.66$^{+0.65}_{-0.49}$&-2.82$^{+0.91}_{-0.71}$\\\\ NGC288/B22  & 1407.0955 &3$\\times$1200& 17 & 3885-6292 & 12130 & 4.01 & 2.48 & -43.8$\\pm$0.9&+0.28$^{+0.29}_{-0.38}$&-2.36$^{+1.26}_{-0.64}$\\\\ NGC288/B186 & 1409.0862 &3$\\times$1200& 31 & 3885-6292 & 11390 & 3.94 & 3.89 & -43.8$\\pm$0.5&+0.17$^{+0.17}_{-0.21}$&-1.45$^{+0.65}_{-0.85}$\\\\ NGC288/B302 & 1408.0891 &3$\\times$1200& 26 & 3885-6292 & 13230 & 4.15 & 1.85 & -43.8$\\pm$0.5&+0.61$^{+0.34}_{-0.32}$&-2.63$^{+0.88}_{-0.66}$\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}    BHB stars with $T_{\\rm eff} \\succeq$ 11,500K show several observational anomalies. By studying the position of BHB stars of globular clusters\\footnote{Most of the known BHB stars are found in globular clusters.} on Hertzspruing-Russell diagrams, it is found that they exhibit photometric jumps (Grundhal et al. 1999) and gaps (Ferraro et al. 1998) at $T_{\\rm eff} \\simeq $11,500~K. These stars also show large abundance anomalies (Glaspey et al. 1989; Behr et al. 1999; Moehler et al. 1999; Behr 2003a). For instance, iron is generally overabundant in these stars relative to their cluster abundance. It is also found that the spectroscopic gravities of these stars are lower than those predicted by the canonical models (Crocker et al. 1988; Moehler et al. 1995). It is commonly believed that atomic diffusion of the elements in the atmosphere of these stars is responsible for the aforementioned observational anomalies. Atomic diffusion arises from the competition between radiative acceleration and gravitational settling. This can produce a net acceleration on atoms and ions, which results in their diffusion \\cite{Michaud70} and which may lead to abundance anomalies.  In addition to the detection of vertically stratified abundances in some BHB stars (Khalack et al. 2007 and 2008, more information is given below), BHB stars with $T_{\\rm eff} \\succeq$ 11,500K have lower rotational velocities than their cooler counterparts (Peterson et al. 1995; Behr et al. 2000a and 2000b; Behr 2003b). These observations suggest that the atmospheres of such stars are stable enough to have atomic diffusion. Theoretical results of Quievy et al. (2009) show that helium sinks in stars with low rotational velocities, which leads to the disappearance of the superficial He convection zone. This then opens the door for atomic diffusion to play a role.  Other theoretical results also concur with the belief that diffusion is present in the atmospheres of BHB stars with $T_{\\rm eff} \\succeq$ 11,500K. For instance, the atmospheric models of Hui-Bon-Hoa, LeBlanc \\& Hauschildt \\shortcite{Hui-Bon-Hoa+00} showed that the observed photometric jumps and gaps for hot BHB stars can be explained by elemental diffusion in their atmosphere. These models calculate self-consistently the structure of the atmosphere while taking into account the stratification predicted by diffusion (assuming equilibrium). These models have recently been improved (LeBlanc et al. 2009) and confirm that vertical stratification of the elements can strongly modify the structure of the atmospheres of BHB stars. Such structural changes for the atmosphere lead to the photometric anomalies discussed above.    As mentioned above, by synthesizing spectral line profiles, Khalack et al. \\shortcite{Khalack+07} found vertical abundance stratification of sulfur in the atmosphere of the field BHB star HD~135485. Vertical stratification of iron abundance was also detected by Khalack et al. \\shortcite{Khalack+08a} in the atmosphere of two BHB stars in the globular cluster M15 and for one star in the globular cluster M13. Also, more recently, Hubrig et al. \\shortcite{Hub09} have found isotopic anomalies of calcium in six BHB stars.  In this paper we will attempt to detect signatures of vertical abundance stratification of iron in a set of fourteen BHB stars from line profile analysis for which we have appropriate data. This extension of the studies of Khalack et al. (2007 and 2008) will serve to provide a larger amount of observational data for iron stratification which can then be used to verify theoretical models. We have selected the BHB stars from the spectra of Behr \\shortcite{Behr03a} relying on their value of effective temperature and their low $V\\sin{i}$ ($\\leq$ 10 km s$^{-1}$) so that atomic diffusion phenomena would most probably be present and detectable in these stars. The properties of the acquired spectra are discussed in Sec.~\\ref{obs}, while in Sec.~\\ref{mod} the details concerning the simulation routine and adopted atmospheric parameters for the program stars are described. The evidence for vertical stratification of iron is given in Sec.~\\ref{vert} along with the estimation of the mean iron abundance, $V\\sin{i}$  and radial velocities for each BHB star considered here. A discussion follows in Sec.~\\ref{discuss}.   \\section[]{Observations} \\label{obs}  The selected BHB stars were observed with the Keck I telescope and the HIRES spectrograph. The stars were selected from the large sample of stars observed by Behr \\shortcite{Behr03a}. The criteria by which the stars considered here were chosen is that their $T_{\\rm eff}$ is close to or larger than 11,500 K and that they possess a low $V\\sin{i}$ ($\\leq$10 km s$^{-1}$). These criteria were chosen to maximize the chance that atomic diffusion is present in their atmosphere and that its effects are potentially observable.  Table~\\ref{tab1} summarizes the information concerning the observations of the fourteen stars chosen by giving (in individual columns) the object identification, the heliocentric Julian Date of the observation, the exposure time, the S/N per pixel the spectral coverage, the size of the seeing disk (FWHM), the effective temperature, the logarithm gravity and the heliocentric radial velocity.  The size of the C1 slit is 0.86 arcsec which provides a spectral resolution of $R=\\lambda/\\delta\\lambda=$45,000. Behr \\shortcite{Behr03a} has found that for the aforementioned stars the underfilling of the slit should not change the estimated spectral resolution by more than 4\\%--7\\%. To process the spectra, Behr~\\shortcite{Behr03a} has employed the package of routines developed by McCarthy \\shortcite{McCarthy90} for the FIGARO data analysis package \\cite{Shortridge93}. A comprehensive description of the data acquisition and reduction procedure is presented by Behr~\\shortcite{Behr03a}.   \\begin{table} \\begin{minipage}{70mm} \\caption[]{Properties of the BHB stars studied here.} \\label{aver} \\begin{tabular}{@{}lcccc@{}} \\hline Cluster/Star& [Fe/H] & $V\\sin{i}$ &$V_{\\rm r}$ & $n$\\\\ &        &(km s$^{-1}$)&(km s$^{-1}$)& \\\\ \\hline M13/WF2-820 & -0.09$\\pm$0.32 & 4.2$\\pm$1.0 &-239.6$\\pm$0.8 & 79\\\\ M13/WF2-2692& +0.03$\\pm$0.25 & 4.4$\\pm$1.0 &-236.2$\\pm$1.0 & 74\\\\ M13/WF2-3123& -0.57$\\pm$0.17 & 6.0$\\pm$1.5 &-238.3$\\pm$1.0 & 21\\\\ M13/WF3-548 & -0.49$\\pm$0.16 & 4.2$\\pm$0.7 &-235.0$\\pm$1.3 & 26\\\\ M13/WF3-1718& +0.14$\\pm$0.13 & 2.0$\\pm$0.3 &-243.7$\\pm$0.7 & 96\\\\ M15/B203    & -0.14$\\pm$0.22 & 5.1$\\pm$1.3 & -95.0$\\pm$0.3 & 32\\\\ M15/B315    & +0.10$\\pm$0.21 & 1.8$\\pm$0.2 &-103.2$\\pm$0.9 & 72\\\\ M15/B334    & -1.38$\\pm$0.86 &10.1$\\pm$0.5 &-106.5$\\pm$2.1 & 11\\\\ M15/B374    & +0.59$\\pm$0.21 & 4.0$\\pm$0.9 &-107.0$\\pm$1.1 & 82\\\\ M92/B176    & -1.26$\\pm$0.75 & 7.0$\\pm$1.0 &-117.7$\\pm$1.2 & 21\\\\ NGC288/B16  & +0.69$\\pm$0.15 & 2.9$\\pm$0.9 & -45.1$\\pm$0.9 & 42\\\\ NGC288/B22  & +0.52$\\pm$0.22 & 2.5$\\pm$0.5 & -43.9$\\pm$0.8 & 47\\\\ NGC288/B186 & +0.40$\\pm$0.17 & 3.3$\\pm$0.8 & -43.9$\\pm$0.5 & 65\\\\ NGC288/B302 & +0.77$\\pm$0.17 & 1.8$\\pm$0.2 & -44.0$\\pm$0.5 & 75\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table}    \\section[]{Simulation procedure} \\label{mod}  The line profile simulations were performed with the {\\sc Zeeman2} spectrum synthesis code (Landstreet \\shortcite{Landstreet88}, Wade et al. \\shortcite{Wade+01}). Khalack \\& Wade \\shortcite{Khalack+Wade06} have modified the {\\sc Zeeman2} code to allow for an automatic minimization of the model parameters using the {\\it downhill simplex method} \\cite{press+}. To analyze the vertical abundance stratification we have built the dependence of the abundance derived from each analyzed profile relative to the optical depth $\\tau_{\\rm 5000}$ for the list of selected line profiles, assuming that the profile is formed mainly at line optical depth $\\tau_{\\rm \\ell}$=1. This depth corresponds to a continuum optical depth $\\tau_{\\rm 5000}$ of a certain value that in turn corresponds to a given layer of the stellar atmosphere model. A comprehensive description of the procedure is given by Khalack et al. \\shortcite{Khalack+07}. This method was also used by Khalack et al. \\shortcite{Khalack+08a} % for the study of stratification in BHB stars and by Thiam et al. \\shortcite{Thiam+10} for HgMn stars.  The stellar atmosphere models used for our analysis were calculated with the {\\sc Phoenix} code \\cite{Hauschildt+97} assuming LTE (Local Thermodynamic Equilibrium), solar metallicity except for the iron abundance and depleted helium abundances, and the surface gravities % extracted from Behr \\shortcite{Behr03a} which are listed in Table~\\ref{tab1}. The depleted helium abundance has also been taken into account during line profile simulations using {\\sc Zeeman2}.     In a previous study, \\cite{Khalack+08a} showed that inclusion of non-zero microturbulence leads to the amplification of signatures of vertical abundance stratification for different chemical species. Therefore, in order to mitigate the influence of microturbulence on the estimation of iron vertical stratification, we have assumed zero microturbulent velocity in the simulations presented here.   In this paper, we have selected a list of % Fe\\,{\\sc ii} lines that are suitable for abundance and stratification analysis. The full list of the analyzed spectral lines is given in the electronic edition of {\\it MNRAS}. Atomic data for Fe\\,{\\sc ii} lines are extracted from VALD-2 (Kupka~et~al.~\\shortcite{Kupka+99}, Ryabchikova~et~al.~\\shortcite{Ryab+99}), Pickering et al. \\shortcite{Pickering+01} and % Raassen \\& Uylings\\footnote{\\rm ftp://ftp.wins.uva.nl/pub/orth} \\shortcite{RU+98}) line database. ",
        "conclusions": "\\label{discuss}  In this paper, we have found additional evidence that iron abundance is vertically stratified in the atmosphere of BHB stars. From the analysis of 14 BHB stars, selected as the most probable candidates for such detections due to their $T_{\\rm eff} \\succeq$ 11,500K and $V\\sin{i} \\leq$ 10 km s$^{-1}$, we have found clear signatures of iron stratification in five stars: B22, B186 in the globular cluster NGC~288; WF2-820, WF2-2692 in M13 and B203 in M15. Two other stars, B334 in M15 and B176 in M92, show extremely strong increase of iron abundance as a function of optical depth in their atmosphere (see Tab.~\\ref{tab3} and Fig.~\\ref{fig4}). However, this result may be influenced by the strongly depleted mean iron abundance (that results in a smaller number of analyzed line profiles). Also, because of other anomalies found for these two stars (see Sect.~\\ref{trac}), some doubts remain about observed stratification of iron in their atmosphere.  For B334 and B176 the effective temperature is only slightly smaller than the limit ($T_{\\rm eff}=$11,500~K) above which BHB stars show the photometric jumps and gaps (Grundahl et al.~\\shortcite{Grundahl+99}, Ferraro et al.~\\shortcite{Ferraro+98}) and where diffusion is efficient. On the other hand, Khalack et al.~\\shortcite{Khalack+08a} have shown that stellar atmosphere models with higher $T_{\\rm eff}$ usually slightly decrease the slope of the element's abundance with respect to the logarithm of excitation potential as well as with respect to $\\log{\\tau_{5000}}$. Therefore, adoption of higher effective temperatures for B334 and B176 will decrease the estimated slopes of iron abundance which increases towards the deeper atmosphere.  The values of effective temperature for the sample of analyzed BHB stars were derived from Kurucz's ATLAS9 synthetic photometry grids \\cite{Behr03a}, which use the stellar atmosphere models with homogeneous distribution of element's abundance. Hui-Bon-Hoa et al.~\\shortcite{Hui-Bon-Hoa+00} and LeBlanc et al. (2010, in preparation) have shown that the observable photometric jumps and gaps can be explained in terms of atomic diffusion mechanism that leads to stratification of element's abundance with atmospheric depths. The models with stratified abundance of chemical species result in different photometric colors as compared to chemically homogenous models. This fact may introduce some errors in the evaluation of $T_{\\rm eff}$ by photometric means when using chemically homogeneous atmospheric models. The uncertainty on $T_{\\rm eff}$ determination should be studied by using the new atmospheric models of LeBlanc et al. (2009) that take into account vertical stratification of the elements. This is outside the scope of the present paper.  Our analysis of the dependency of the available data for vertical iron stratification in BHB stars with respect to the effective temperature results in the detection of a trend (see Fig.~\\ref{fig4}). The same results were obtained from the modeling of iron vertical stratification employing the self-consistent stellar atmosphere models with vertically stratified abundance of chemical species (LeBlanc et al. 2010, in preparation). The results presented here show that for BHB stars with $T_{\\rm eff}$ around 11,500K have the strongest vertical iron stratification gradients. This gradient diminishes as $T_{\\rm eff}$ increases up to $T_{\\rm eff} \\simeq$ 14,000K, where no significant iron stratification exists. Nevertheless, to statistically improve the registered trend of iron stratification more data for BHB stars with appropriate values of $T_{\\rm eff}$ and $V\\sin{i}$, especially for $T_{\\rm eff}$ near 11,500K need to be analyzed. Such additional data will be useful to constrain theoretical model atmospheres that include stratification of elements, similar to those of LeBlanc et al. (2009).   \\begin{figure} \\includegraphics[scale=0.32,angle=-90]{BHB-Teff-Strat53.eps} \\caption{Plot of the slope of iron vertical stratification versus effective temperature of analyzed BHB stars (see text for details). The dashed line is a linear fit that includes all stars except for the two stars represented by empty circles (B176 and B334). The linear fit given by the solid line also excludes the star WF3-1718 represented by the empty square.} \\label{fig4} \\end{figure}"
    },
    "1005/1005.2659_arXiv.txt": {
        "abstract": "{The stars in a  Globular cluster (GC)  have always been considered coeval and of the same metallicity. Recently, this assumption has been challenged on the basis of  spectroscopic and photometric observations, which show the existence of various generations of stars in GCs, differing in the abundances of products of H-burning at high temperatures. The main final product of this burning is He. It is then important to study the connections between stars properties and He content.  We consider here the about 1400 stars on the Red Giant Branch (RGB) observed with FLAMES@VLT in 19 Galactic Globular Clusters  (GCs) in the course of our project on the Na-O anticorrelation. Stars with different He are expected to have different temperatures (i.e., different colours),  slightly different metallicities [Fe/H], and different luminosity levels of the RGB bump. All these differences are small, but our study has the necessary precision, good statistics, and homogeneity to detect them. Besides considering the observed colours and the temperatures and metallicities determined in our survey,  we computed suitable sets of   stellar models - fully consistent with those present in the BaSTI   archive - for various assumptions about the initial helium content. We find that differences in observable quantities that can be attributed to  variations in He content are generally detectable between stars of the Primordial (P, first-generation) and Extreme (E, second-generation) populations, but not between the Primordial and Intermediate ones (I).  The only exception, where  differences are significant also between P and I populations, is the cluster NGC~2808, where three populations are clearly separated also on the Main Sequence and possibly on the Horizontal Branch. The average enhancement in the He mass fraction Y between P and E stars is about 0.05-0.11, depending on the assumptions.   The differences in Y, for NGC~2808 alone, are of about 0.11-0.14 between  P and I stars, and about 0.15-0.19 between P and E stars, again depending on the assumptions.  When we consider the RGB bump luminosity of first and second-generation  stars we find different levels; the implied   Y difference is more difficult to quantify, but is in agreement with the other determinations. } ",
        "introduction": "Recent progresses have indicated that large star-to-star variations exist in the He content among stars in globular clusters (GCs), in spite of difficulties in direct derivation of He abundances  \\citep[see however][]{villanova09,moehler07}, The first clear evidence has been obtained from observations of multiple main sequences in \\object{$\\omega$~Cen}  \\citep{bedin04} and NGC~2808 \\citep{dantona05,piotto07}. In $\\omega$~Cen, stars on the bluest main sequence (MS) are more metal-rich than those of the reddest one \\citep{piotto05}, a fact that can only be explained by an He content higher by about 0.10 to 0.15 in the mass fraction Y \\citep[as originally suggested by][]{norris04} with respect to a canonical (primordial) He content of about 0.245 commonly adopted for the bulk of the stellar population in this cluster. The small scatter found in NGC~2808 among red giant branch stars also implies a larger He content for the bluer MS, by roughly the same amount. Two other clusters have been found to have wide MSs: 47 Tuc \\citep{anderson09} and  NGC~6752 \\citep[][who also suggest a possible split]{milone10}. \\footnote{The split subgiant branches (SGB, e.g., NGC~6388: \\citealt{moretti09}; \\object{NGC~1851}: \\citealt{milone08}) seem more related to differences in age or CNO content. A spread in age is probably the cause of the spread/split MS turn-off's found in many intermediate-age Magellanic Clouds GCs (\\citealt{mackey,milone09} in LMC;  \\citealt{glatt} in SMC) even if for them -but notably not for the old Galactic GCs-  there is an alternative physical interpretation based on the observational effects associated to the occurrence of stellar rotation as discussed by \\cite{selma}. }  A star-to-star spread in the He abundance may explain many aspects of the horizontal branches (HB) of GCs, as extensively discussed in \\cite{gratton10}, but we are aware that also the case for a null spread of He in GCs is maintained in literature \\citep[e.g.,][]{catelan}. Internal  variations in He are likely connected with other chemical signatures found in GCs, related to H burning occurring at high temperature. This has been first proposed in a modern way by \\cite{dantona02}, who noticed that the smaller masses of He-rich stars may naturally explain the correlation between the presence of stars at the blue extreme of the horizontal branch, and of stars extremely depleted in O in various clusters. \\cite{lee05} suggested that a super-He rich population could explain the extreme HB stars in GCs, specifically presenting the cases of  $\\omega$ Cen and NGC~2808 (where they predicted a spread or a split of the MS, just before it was actually found). \\cite{kaviraj} suggested that the far-UV properties of GCs in the Virgo elliptical M87 could be explained by a super-He rich population similar to the one in $\\omega$ Cen. \\cite{dac08}  tried to explain the whole distribution of stars along the HB of various clusters (including NGC~2808, M~13, and M~3) as due to star-to-star variations in the He content (but see also \\citealt{castellani05}). In a companion paper \\citep{gratton10}, we make a more systematic comparison of the extremes of the distribution of stars along the HB and of the Na-O anticorrelation. We find that median HB colours are quite well reproduced by a combination of metallicity, age, and a simple mass loss proportional to metallicity. However, the low-mass extreme of the HB is clearly correlated with the extension of the Na-O anticorrelation \\citep[see also][]{carretta09a} which, in turn, seems to be linked with the cluster absolute magnitude, a proxy for cluster mass.  These correlations suggest that the multiple generation scenario, needed to explain the Na-O anticorrelation,\\footnote{It is generally assumed that first-generation stars have chemical composition similar to field stars of similar age and metallicity (i.e., high O and Mg, low Na and Al, etc.), while second-generation stars, formed from material polluted by matter lost from the primordial population, show the abundance patterns (e.g., Na-O anticorrelation) that are unique of GCs. } is able to explain many peculiarities of GCs. As we have seen, a crucial datum in this scenario is the He content of second-generation stars. While the most spectacular evidence of a variation of the He content comes from the HB and the MS, also the red giant branch (RGB) stars are very useful because chemical composition (but -alas- not He content) can be determined easily enough for these stars using spectroscopy, unlike faint MS stars. This allows us to cross correlate the abundances of various elements, providing important constraints on the properties of the He producer, which have not been clearly identified  yet (two main possible mechanisms have been proposed, intermediate/massive asymptotic giant stars -AGB- and fast rotating massive stars  -FRMS- see e.g., \\citealt{dav07} and \\citealt{decressin},  respectively). Various properties of RGB stars are affected by the variable He abundances: \\begin{enumerate} \\item The temperature (at a given luminosity) of He-rich RGB stars is slightly warmer than that of He-poor ones \\citep[see e.g.,][]{dantona02}. The difference is not large (only a few tens of K), and very accurate relative determinations are required. However, both colours and line excitation can be used to show that this really occurs. \\item The gravity (again at a given luminosity) is also different, due to both the smaller mass and higher temperature of He-rich stars. He-rich stars have slightly lower gravity, but this effect is very small, and difficult to show up, because of several complications (for instance, the larger molecular weight and reduced continuum opacity simulate a higher gravity, offsetting most of the effect expected on e.g., ionisation equilibrium) so we will not consider it here. \\item The [Fe/H] value is different, because of the different H content (broadly speaking, if Y increases, X decreases, hence [Fe/H] increases; things are of course more complicated, see Sec.~\\ref{deltamet}): this difference is small but not negligible. We already noticed this effect in NGC~2808 \\citep[][Paper I]{carretta06}. \\item Finally, the luminosity of the RGB bump is also different: the RGB bump is more luminous for He-rich stars \\citep{iben,salaris06}, even if also in this case the situation is more complex (see Sect.~\\ref{bump}). We showed that possible evidence for this effect can be traced in the run of Na abundances along the RGB of NGC~6218 (=M~12) and NGC~6752 \\citep[][Paper IV]{carretta07b}. \\end{enumerate}  In this paper, we intend to look for evidence of a variable He content in the sample of more than 1,400 stars along the RGB of 19 GCs for which we acquired FLAMES/GIRAFFE spectra in the last few years \\citep[presented in a series of  papers: see][-Paper VII and VIII- and references therein]{carretta09a,carretta09b}. Several features of this data set make it very promising for the present purpose: the sample is very large; it has been analysed using a very homogenous technique; star-to-star errors in effective temperatures, usually the main source of errors in this analysis, were reduced to very small values by a carefully tailored technique. In addition, we have determined Na and O abundances for almost all these stars, allowing to classify them into three different groups according to Na enhancement and O depletion (Paper VII). We called these groups P=primordial, that is stars with composition similar to that of field stars, and likely belonging to the primordial population; I=intermediate, and E=extreme, that are stars with different degrees of O depletion and Na (and likely also He) enhancement. We recall here that we conservatively applied this definition only to stars with both O and Na measured, thus reducing our sample to about 960 objects, the ones that will be used in the present paper. On the other hand, the way we analysed stars should be carefully considered when we try to extract information from our data set. The most significant assumptions we made are that stars of the same luminosity have the same effective temperature and surface gravity. Whenever needed, we will consider explicitly the consequences of these assumptions, and we will correct our results accordingly.  Beside observations, determination of the impact of different He abundances requires appropriate modelling. To this purpose we used a set of BaSTI evolutionary  models \\citep{pietrinferni04,pietrinferni06,pietrinferni09} purposely computed for this paper employing two sets of heavy elements mixtures (see Sect.~\\ref{sec2} and \\ref{bump}  for details and relations).  The structure of this paper is as follows: in Sec.~\\ref{sec2} we present the observed differences in colour, temperature, and metallicity and derive the implied differences in He content; in Sec.~\\ref{bump} we compare the RGB bumps of first and second-generation stars, deducing a different He abundance also in this case. Finally, we summarise and discuss our results in Sec.~\\ref{ultima}. ",
        "conclusions": "\\label{ultima}  Recently, spreads and splits in GC RGBs other than the ones  (in $V-K$, metallicity and temperature) considered here have been presented. They are most probably due to the same phenomenon we are seeing here, since the He-enhanced stars should also show other chemical signatures. For instance, the correlated N-enhancement should be responsible for the effects seen in the colour-magnitude diagrams involving the Johnson $U$ or Str\\\" omgren $u$ filters \\citep{marino08,yong08}. On the other hand, the effect seen in  $hk$ photometry, attributed to a  spread in calcium by \\cite{lee09}, and demonstrated not to be so by \\cite{carretta10}, still awaits a definite explanation.  We have used information on RGB stars in GCs to infer the plausible He differences implied by the existence of two generations of stars. This is a new approach, since the different methods used in the past to deduce the He content of GCs generally involve the HB,  the evolutionary phase more sensitive to even small He variations (the two exceptions being $\\omega$~Cen  and NGC~2808 with their multiple MSs).  Apart from the information coming from photometric data, He can be deduced from spectra of HB stars, where He lines are visible for temperatures hotter than about 8500 K. \\cite{villanova09}  studied seven HB stars in NGC~6752, using UVES spectra at high-resolution and very high S/N. For the four stars where they were able to measure the He line, they obtained a value similar to the cosmological one (Y=0.245). However, these four are all O-rich, Na-poor, while  Villanova et al.  could not measure He for the only Na-rich, O-poor star, the one that should  also be He-enhanced, so the case is not yet settled. However, the photospheric abundances of HB stars hotter than about 11500 K are altered  by atomic diffusion processes and cannot be used directly to measure the original values.  For instance, Behr (2003) found a depletion in moderately hot HB stars in several GCs. Moehler et al., in a series of papers \\citep[e.g.,][]{moehler07} concentrated instead on the hottest part of the HB (and blue-hook stars in the two GCs -$\\omega$ Cen and NGC~2808- where they are present) and found He-enhancement. This can be attributed either to  pollution from a previous generation  or to  He flash induced mixing occurring in hot He-flashers, i.e. those stars experiencing the He flash not at the RGB tip but along the White Dwarf cooling sequence \\citep[see][]{cc93}.  A further discussion of this problem in $\\omega$ Cen can be found in \\cite{cassisi09}.  We have studied in a companion paper \\citep{gratton10} the effect of (even a small) difference in He content on the morphology of HBs, and found that He is most probably the third parameter governing the HB, after metallicity and age. We refer to that paper for a long and detailed discussion on the effects of He variations, on methods to measure He content \\citep[e.g., the R-method:][]{iben}, on comparison between results obtained through different approaches. In the present paper we limit our analysis to the differences in He that can be deduced from RGB stars.  May we use the present analysis, combined with the previous work on the Na-O and other (anti)correlations to try discriminating between  AGB and FRMS as the polluters of the second-generation stars? We need to consider that He is produced essentially in MS for both  classes of polluters, while there is a difference for the other ``peculiar\" elements (O, Na, Mg, Al, Si), produced during  Hot Bottom Burning or in MS, respectively. However, we still miss important details of  stellar evolution (e.g., how much He is dredged up in AGB stars, or what is the dependence on metallicity -via rotation- of FRMS), the numbers are still small and the inferences are not  conclusive. We need to further investigate the dependencies of variations of He and  those other elements  on the stars and cluster properties.  In summary, with the present work we have seen that: \\begin{itemize} \\item[i)] it is possible to deduce variations of He also from RGB stars in GCs when using large samples, treated homogeneously; \\item[ii)] these variations are generally measurable from colour (here $V-K$) and [Fe/H] only between first-generation (P) and extreme second-generations (E) stars; \\item[iii)] the variations implied for the average of the 19 GCs considered are of the order of 0.05-0.10 ($\\Delta$Y); however, the heavy-elements mixture assumed for the second-generation stars has to be taken into account (and we have presented two extreme cases); \\item[iv)] NGC~2808 is a notable exception: in this cluster, differences in colour and metallicity are seen also between P and the intermediate (I) component of second-generation stars. This is perfectly in line with the fact that NGC~2808 presents a clear evidence of three He levels also in MS. For NGC~2808 the deduced $\\Delta$Y values are higher (see Table~\\ref{t:deltaY}) than for the average of all GCs; \\item[v)] similar results can be obtained also considering the luminosity of the RGB bump.  We could test this method only for 14 GCs; unluckily, NGC~2808 is not among them. Within the limits of the rather poor statistics,  we found that the implied $\\Delta$Y values are in reasonable agreement with those found from colour and metallicity; \\item[vi)] the absolute calibration of the $\\Delta$Y in GCs is still to be defined. \\end{itemize}"
    },
    "1005/1005.0006_arXiv.txt": {
        "abstract": "We show that the mass-metallicity relation observed in the local universe is due to a more general relation between stellar mass \\mstar, gas-phase metallicity and star formation rate (SFR). Local galaxies define a tight surface in this 3D space, the Fundamental Metallicity Relation (FMR), with a small residual dispersion of $\\sim$0.05~dex in metallicity, i.e, $\\sim$12\\%. At low stellar mass, metallicity decreases sharply with increasing SFR, while at high stellar mass, metallicity does not depend on SFR.  High redshift galaxies, up to z$\\sim$2.5 are found to follow the same FMR defined by local SDSS galaxies, with no indication of evolution. In this respect, the FMR defines the properties of metal enrichment of galaxies in the last 80\\% of cosmic time. The evolution of the mass-metallicity relation observed up to z=2.5 is due to the fact that galaxies with progressively higher SFRs, and therefore lower metallicities, are selected at increasing redshifts, sampling different parts of the same FMR.  By introducing the new quantity $\\mu_\\alpha=\\rm{log}(M_\\star)-\\alpha\\,\\rm{log}(SFR)$, with $\\alpha$=0.32, we define a projection of the FMR that minimizes the metallicity scatter of local galaxies.  The same quantity also cancels out any redshift evolution up to z$\\sim$2.5, i.e, all galaxies follow the same relation between $\\mu_{0.32}$ and metallicity and have the same range of values of $\\mu_{0.32}$. At z$>$2.5, evolution of about 0.6~dex off the FMR is observed, with high-redshift galaxies showing lower metallicities.  The existence of the FMR can be explained by the interplay of infall of pristine gas and outflow of enriched material. The former effect is responsible for the dependence of metallicity with SFR and is the dominant effect at high-redshift, while the latter introduces the dependence on stellar mass and dominates at low redshift. The combination of these two effects, together with the Schmidt-Kennicutt law, explains the shape of the FMR and the role of $\\mu_{0.32}$. The small metallicity scatter around the FMR supports the smooth infall scenario of gas accretion in the local universe. ",
        "introduction": "\\label{sec:intro}  Gas metallicity is regulated by a complex interplay between star formation, infall of metal-poor gas and outflow of enriched material. A relation between magnitude and metallicity was discovered in the '70 \\citep{McClure68,Lequeux79}, in which more luminous galaxies also have higher metallicities. Later on, it was understood that this luminosity-metallicity relation is a manifestation of a more fundamental stellar mass-metallicity relation where galaxies with larger stellar mass \\mstar\\ have higher metallicities \\citep{Garnett02,Perez-Gonzalez03a,Pilyugin04,Tremonti04,Savaglio05,Lee06, Cowie08,Panter08,Kewley08,Hayashi09a,Michel-Dansac08,Liu08,Rodrigues08, Perez-Montero09}.  The origin of this relation is debated, and many different explanations have been proposed, including ejection of metal-enriched gas (e.g., \\citealt{Edmunds90,Lehnert96a,Garnett02,Tremonti04,Kobayashi07,Scannapieco08,Spitoni10}), ``downsizing'', i.e., a systematic dependence of the efficiency of star formation with galaxy mass (e.g., \\citealt{Brooks07,Mouchine08,Calura09a}), variation of the IMF with galaxy mass \\citep{Koppen07}, and infall of metal-poor gas \\citep{Finlator08,Dave10}.  Recently, evidence has been found that at high redshift the infall of pristine gas can have a dominant role \\citep{Bournaud09,Dekel09,Brooks09,Agertz09b}. The mass-metallicity relation has been studied by \\cite{Erb06a}  at z$\\sim$2.2 and by \\cite{Maiolino08} and \\cite{Mannucci09b} at z=3--4, finding a strong and monotonic evolution, with metallicity decreasing with redshift at a given mass. The same authors \\citep{Erb06a,Erb08,Mannucci09b} have also studied the relation between metallicity and gas fraction, i.e., the effective yields, obtaining clear evidence of the presence of infall in high redshift galaxies.  Even more recently, \\cite{Cresci10} have obtained the evidence of ``positive'' metallicity gradients in a small sample of three disk galaxies at z=3.5, with lower metallicities in more active regions of star formation. The presence of these gradients in galaxies with very regular dynamics can be understood as a consequence of accretion of metal-poor gas, producing a new episode of star- formation in older, more metal-rich galaxies.  If infall is at the origin of the star formation activity, and outflows are produced by exploding supernovae (SNe), a relation between metallicity and SFR is likely to exist. In other words, SFR is a parameter that should be considered in the scaling relations that include metallicity. The role of specific SFR (SSFR) in this content was already studied by \\cite{Ellison08a}, who presented a mild ($\\le$0.1~dex) dependence of metallicity on SSFR when binning galaxies according to their \\mstar. Recently, also \\cite{Lopez-Sanchez10c} presented evidence for a link between SFR and metallicity, while \\cite{Peeples09} reported high values of SFR in a sample of outliers, toward low metallicities, of the mass-metallicity relation.  To test the hypothesis of a correlation between SFR and metallicity in the present universe and at high redshift, we have studied several samples of galaxies at different redshifts whose metallicity, \\mstar, and SFR have been measured. In the next section we present the data samples we are using for this study. In sec.~\\ref{sec:massmet} we study the mass-metallicity relation as a function of SFR, and in the following section we introduce the Fundamental Metallicity Relation. In sec.~\\ref{sec:discussion} we discuss the physical origin of this relation. We adopt a $\\Lambda$CDM cosmology with $H_0$=70 km/sec, $\\Omega_m$=0.3 and $\\Omega_\\Lambda$=0.7. Stellar mass \\mstar\\ and SFR are expressed in \\msun\\ and in \\msun/yr, respectively.\\\\   \\begin{figure*} \\centerline{ \\includegraphics[width=0.48\\textwidth]{massmet.pdf} \\includegraphics[width=0.48\\textwidth]{sfrmet.pdf} } \\caption{ {\\em Left panel:} The mass-metallicity relation of local SDSS galaxies. The grey-shaded areas contain 64\\% and 90\\% of all SDSS galaxies, with the thick central line showing the median relation. The colored lines show the median metallicities, as a function of \\mstar, of SDSS galaxies with different values of SFR. {\\em Right panel:} median metallicity as a function of SFR for galaxies of different \\mstar. At all \\mstar\\ with log(\\mstar)$<$10.7, metallicity decreases with increasing SFR at constant mass . } \\label{fig:massmet} \\end{figure*} ",
        "conclusions": "We have studied the dependence of gas-phase metallicity 12+log(O/H) on stellar mass \\mstar\\ and SFR on a few samples of galaxies from z=0 to z=3.3. In the local universe, we find that metallicity is tightly related to both \\mstar\\ and SFR (fig.~\\ref{fig:massmet}), and this defines the Fundamental Metallicity Relation (fig.~\\ref{fig:cfr1}). The residual metallicity dispersion of local SDSS galaxies around this FMR is about 0.05~dex (fig.~\\ref{fig:plotdif}), i.e, about 12\\%. The well-known mass-metallicity relation, together with the luminosity-metallicity and velocity-metallicity relations, is one particular projection of this relation into one plane, and neglecting the dependence of metallicity on SFR results in doubling the observed dispersion.   When high redshift galaxies are compared to the FMR defined locally, we find no evolution up to z=2.5, i.e., high-redshift galaxies follow the same FMR defined by SDSS galaxies even if they have higher SFRs (fig~\\ref{fig:plotevol}). This means that the same physical processes are in place in the local universe and at high redshifts. The observed evolution of the mass-metallicity relation is due to the increase of the average SFR with redshift, which results in sampling different parts of the same FMR at different redshifts.  At even higher redshift, z$\\sim$3.3, evolution of $\\sim$0.6~dex with respect to the FMR is found, although several observational effects and selection biases may affect the size of this evolution. This is an indication that different mechanisms start to dominate.  Even if the nature of the FMR is 3-dimensional in \\mstar, SFR and 12+log(O/H), metallicity is found to be tightly correlated with $\\mu_{0.32}=\\rm{log}(M_\\star)-0.32\\,\\rm{log}(\\rm{SFR})$. Galaxies at any redshifts up to z=2.5 follow the same $\\mu_{0.32}$-metallicity relation and have the same range of values of $\\mu_{0.32}$ (fig.\\ref{fig:bestfit}).  Metallicity in galaxies of any mass is found to have the same dependence of SSFR, with galaxies above the threshold of SSFR=$10^{-10}$yr$^{-1}$  showing a rapidly decreasing metallicities with increasing SSFR (fig.~\\ref{fig:ssfrmet}).  The interpretation of the existence of the FMR, its dependence of SSFR, and the role of $\\mu_{0.32}$ depend on the relevant timescales. If dynamical times are shorter than timescales for chemical enrichment, the dependence of metallicity on SFR can be easily explained by dilution by infall. In this case this effect dominates the metallicity evolution of galaxies at high redshift, when galaxies grow because of massive accretions of metal poor gas and produce large SFR. Also galaxies at z=3.3 can fit into this scheme, with large masses of infalling gas. This is in agreement with other recent independent results \\citep{Mannucci09b,Cresci10}. In the local universe, galaxies with large SFR are rare and often associated to merging events, and other effects becomes dominant which relate metallicity mainly to \\mstar. Outflow is an possibly, although downsizing could also work. If, in contrast, infall and SFR evolve on timescales much longer than the chemical enrichment timescale, a sort of steady-state situation is created: continuos infall of metal-poor gas, which both sustains SFR and dilutes metallicity, and outflow of metal-rich gas in galactic winds. In this case the outflows must depend on both  mass and SFR.  The small residual scatter around the FMR in the local universe supports the smooth accretion scenario, where galaxy grow is dominated by continuos accretion of cold gas. Interacting and merging galaxies are expected to show larger spread around the FMR, in agreement to what is actually observed. Galaxies at intermediate and high redshifts show larger metallicity dispersions, which could be due either to uncertainties in the measurements, or to intrinsic dispersion, or both. This effect prevents us to study the evolution of residual scatter with redshift. Nevertheless, the absence of significant biasses in metallicity or in SFR up to z=2.5 points toward the existence of the same physical effects and the dominance of smooth accretion even at intermediate redshift.\\\\  \\noindent {\\bf Acknowledgements}  We are grateful to the MPA/JHU teams which made their measured quantities on SDSS galaxies publicly available, and to the members of the italian Virtual Observatory for support on data visualization. We also acknowledge stimulating discussions with T. Nagao, R. Dav\\'e, S. Ellison, and the Arcetri MEGA group."
    },
    "1005/1005.3719_arXiv.txt": {
        "abstract": "{} {We search for frozen water and its processing around young stellar objects (YSOs of class I/II). We try to detect potential, regional differences in water ice evolution within YSOs, which is relevant to understanding the chemical structure of the progenitors of protoplanetary systems and the evolution of solid materials. Water plays an important role as a reaction bed for rich chemistry and is an indispensable requirement for life as known on Earth.} {We present our analysis of NAOS-CONICA/VLT spectroscopy of water ice at 3~$\\mu$m for the T\\,Tauri star YLW\\,16\\,A in the $\\rho$~Ophiuchi molecular cloud. We obtained spectra for different regions of the circumstellar environment. The observed absorption profiles are deconvolved with the mass extinction profiles of amorphous and crystallized ice measured in laboratory. We take into account both absorption and scattering by ice grains.} {Water ice in YLW16A is detected with optical depths of between $\\tau=1.8$ and $\\tau=2.5$. The profiles that are measured can be fitted predominantly by the extinction profiles of small grains (0.1~$\\mu$m -- 0.3~$\\mu$m) with a small contribution from large grains ($<$$10$\\%). However, an unambiguous trace of grain growth cannot be found. We detected crystallized water ice spectra that have their origin in different regions of the circumstellar environment of the T\\,Tauri star YLW\\,16\\,A. The crystallinity increases in the upper layers of the circumstellar disk, while only amorphous grains exist in the bipolar envelope. As in studies of silicate grains in T\\,Tauri objects, the higher crystallinity in the upper layers of the outer disk regions implies that water ice crystallizes and remains crystallized close to the disk atmosphere where water ice is shielded against hard irradiation. } {} ",
        "introduction": "\\label{section:introduction} The importance of water (ice) to protoplanetary systems is justified by the following properties: liquid water is a proper solvent for many minerals and organic molecules and allows the formation of complex organic molecules. Water is the reaction bed for the photochemical synthesis of both aromatic compounds (Bernstein et al.~\\cite{bernstein}) and amino acids (Munoz Caro et al.~\\cite{munoz}), even in its solid state. Therefore, rich deposits of water (ice) in protoplanetary disks are assumed to be a requirement for the formation of life as known on Earth. The importance of water ice is in addition emphasized by the assumption that ice planets, such as Neptune, form outside the snowline where water condenses on dust grains. The mass fraction of solid matter in protoplanetary disks abruptly increases at the snowline at least by one order of magnitude (Stevenson \\& Lunine~\\cite{stevenson}).  The water ice band at $3.08\\,\\mathrm{\\mu m}$, which is caused by a vibrational excitation of the OH bond in the water molecule, was primarily discovered in the envelopes of numerous deeply embedded protostars (Gillett \\& Forrest~\\cite{gillett}; Gillett et al.~\\cite{gillettII}; Merrill et al.~\\cite{merrill}). Frozen water that can be found in dust particles is the most frequent ice molecule in YSOs (Ehrenfreund et al.~\\cite{ehrenfreund})\\footnote{Fourteen different ice compounds, such as CO, CO$_\\mathrm{2}$, or CH$_\\mathrm{4}$ ice are known. Ice molecules are weakly bounded just by van-der-Waals forces or by hydrogen bonds. A chemical compound that can be found in metals, minerals, and diamonds does not form.}. The huge water ice deposit results from its molecular property: in contrast to nonpolar, volatile ice molecules, such as CO, CO$_\\mathrm{2}$, N$_\\mathrm{2}$, and O$_\\mathrm{2}$ ice, H$_\\mathrm{2}$O ice has a high sublimation temperature of $\\sim$$120\\,\\mathrm{K}$ depending on the ambient pressure (Davis~\\cite{davis}). The polarity of the water molecule is responsible for its stable adsorption by condensation seeds such as silicate dust grains. Cosmic irradiation and chemical reactions could force the evaporation of water ice (L\\'eger et al.~\\cite{leger}; Hartquist \\& Wolf~\\cite{hartquist}). However, these reactions are impeded by regions with strong shielding and low temperatures (Tielens \\& Hagen~\\cite{tielens}). This is why the water ice band is commonly observed in the spectra of deeply embedded protostars (Boogert et al.~\\cite{boogert}). However, it is assumed that larger deposits of water ice exist in the optically thick regions of circumstellar disks (Fig.~1 in Chiang et al.~\\cite{chiang}). This assumption may be supported by the detection of the water ice band at $46.5\\,\\mathrm{\\mu m}$ (Creech-Eakman et al.~\\cite{creech-eakman}) and $3.08\\,\\mathrm{\\mu m}$ (Leinert et al.~\\cite{leinert}; Terada et al.~\\cite{terada}) in the spectra of several T\\,Tauri objects.  The water ice band at $3.08\\,\\mathrm{\\mu m}$ in L band\\footnote{The L band covers a wavelength interval between $2.8\\,\\mathrm{\\mu m}$ and $4.2\\,\\mathrm{\\mu m}$ in the near-infrared range.} is assumed to be an excellent indicator of the evolutionary status of circumstellar disks (van de Bult~\\cite{bult}). In a protostellar molecular cloud, water ice has a primarily amorphous structure, as many previous studies have shown (e.g., Thi et al.~\\cite{thi}; Pontoppidan et al.~\\cite{pontoppidan}). As soon as the temperature exceeds a level of $\\sim$$80\\,\\mathrm{K}$, water ice begins to crystallize (Hagen et al.~\\cite{hagen}). This conversion can be detected by monitoring the profile of the optical depth, i.\\/e., the logarithm of the absorption band: as the crystallinity increases, the profile narrows and its maximum shifts to longer wavelengths (e.g., Smith et al.~\\cite{smith}). But the crystalline structure can also be destroyed, e.g., by hard irradiation (Kouchi \\& Kuroda~\\cite{kouchi}). In addition to crystallization, grain growth of ice particles can be inferred from a modification of the optical depth profile. An increase in the dust grain radius results in a higher scattering efficiency, a broadening of the absorption band at long wavelengths, and a shift of the maximum to longer wavelengths. These changes are comparable to those related to the prominent $10\\,\\mathrm{\\mu m}$ silicate band (e.g., Bouwman et al.~\\cite{bouwman}), but the origin of the spectral bands is different. While the emission band of silicate has its origin in the warm disk atmosphere, the absorption band of water ice originates in more embedded, cold regions closer to the disk midplane and far from the central star.  The water ice band interferes with absorption bands of additional materials that are dissolved in water ice or form a complex compound with the water molecule: ammonia hydrate H$_\\mathrm{2}$O.NH$_\\mathrm{3}$ (minimum at $3.47\\,\\mathrm{\\mu m}$: Merrill et al.~\\cite{merrill}; Mukai et al.~\\cite{mukai}; Dartois \\& d'Hendecourt~\\cite{dartois}), ammonia NH$_\\mathrm{3}$ ($2.96\\,\\mathrm{\\mu m}$: d'Hendecourt et al.~\\cite{dhendecourt}; Dartois \\& d'Hendecourt~\\cite{dhendecourt}), methanol CH$_\\mathrm{3}$OH ($3.54\\,\\mathrm{\\mu m}$: Dartois et al.~\\cite{dartois}), hydrocarbons with the vibrational excitations of the CH, CH$_\\mathrm{2}$, and CH$_\\mathrm{3}$ bond ($3.38\\,\\mathrm{\\mu m}$, $3.42\\,\\mathrm{\\mu m}$, and $3.48\\,\\mathrm{\\mu m}$: Duley \\& Williams~\\cite{duley}; Chiar et al.~\\cite{chiar}), and other chemical compounds (Ehrenfreund et al.~\\cite{ehrenfreund}). The absorption bands of these molecules are narrow and interfere predominantly with the water ice band at longer wavelengths around $3.6\\,\\mathrm{\\mu m}$. They can therefore be distinguished from the water ice band. To study the chemical evaluation of ices in solar-mass systems, a survey in the infrared wavelength range where $41$ low-luminosity YSOs were included was conducted by using the infrared spectrograph on the Spitzer Space Telescope (Boogert et al.~\\cite{boogertII}).  The scientific scope of this study is to look for spatial variations in the water ice band profile at $3.08\\,\\mathrm{\\mu m}$ within a T\\,Tauri star to detect possible, regional differences in water ice evolution. In Sect.~\\ref{section:observation} and \\ref{section:datareduction}, we scrutinize in detail our spectroscopic observations and data reduction, respectively. The derived water ice profiles are presented in Sect.~\\ref{section:results}. The ice profiles that could be derived are analyzed in Sect.~\\ref{section:modelresults} by considering both the absorption profiles of ice grains of different size and crystallization degree that are measured in laboratory. This well-established approach of deconvolution is described in Sect.~\\ref{section:approach}. The question of whether evolved, i.\\/e., large and/or crystyllized ice grains may finally be found is discussed in Sect.~\\ref{section:graingrowth} and \\ref{section:crystalline?}. We conclude this study with a summary in Sect.~\\ref{section:summary}. ",
        "conclusions": "\\label{section:summary} We have presented an L band spectroscopic observation performed with NAOS-CONICA of a YSO that is strongly inclined. These systems allow the study of optical thick circumstellar regions that shield water ice from hard irradiation and evaporation. The absorption feature of water ice at $3.08\\,\\mathrm{\\mu m}$ was identified in the spectra with optical depths of between $1.8$ and $2.5$, depending on the region observed. The optical depth, the water-ice column density, and the visual extinction are higher in the more peripheral regions at greater distance from the photocenter of our observations. The optical depths were derived from the spectra by assuming that a straight line represents the continuum flux.  Considering the extinction profiles of amorphous and crystallized water ice grains with sizes ranging from $0.1\\,\\mathrm{\\mu m}$ up to $1.5\\,\\mathrm{\\mu m}$, the derived optical depths were deconvolved using a fitting routine that was already presented in Schegerer et al.~(\\cite{schegerer}). The optical depths of the spatially unresolved spectra were found to be dominated by the extinction profile of small, amorphous (i.\\/e., non-evolved) water ice grains. However, in spectra for which the T\\,Tauri object YLW\\,16\\,A could be spatially resolved, crystallized water ice could be found in possible outer disk regions, i.\\/e., in disk layers at radii $>$$30\\,\\mathrm{AU}$ in projection. The putative bipolar envelope of this source was not found to exhibit any evidence of crystallized material (Fig.~\\ref{figure:waterice-sketch}). A growth in the size of the water ice grains cannot unambiguously be determined because the corresponding modification of the water ice feature may also be either caused by an increase in the porosity and/or a growth of the particle core.  \\begin{figure*}[!t] \\centering \\resizebox{12.cm}{!}{\\includegraphics[angle=-90]{figure/004fig44.ps}} \\caption{Sketch for a strongly inclined circumstellar disk such as YLW\\,16\\,A. We assume that amorphous, non-evolved water ice is located in the circumstellar disk and in the bipolar envelope. The assumed location of crystallized water ice in the disk is close to upper disk layers, at disk radii that are larger than that of the snowline. The observer and the spectroscopic slit are also sketched. The orthogonal slit orientation is shown. Radiation from amorphous and crystallized ice grains are represented by dashed and solid arrows, respectively.} \\label{figure:waterice-sketch} \\end{figure*}"
    },
    "1005/1005.3834_arXiv.txt": {
        "abstract": "% Traditionally, solar physicists become anxious around solar minimum, as they await the high-latitude sunspot groups of the new cycle.  Now, we are in an extended sunspot minimum with conditions not seen in recent memory, and interest in the sunspot cycle has increased again.  In this paper, I will describe some of the characteristics of the current solar minimum, including its great depth, its extended duration, its weak polar magnetic fields, and its small amount of open flux.  Flux-transport simulations suggest that these characteristics are a consequence of temporal variations of the Sun's large-scale meridional circulation. ",
        "introduction": "When I was asked to give this talk, I wondered if this minimum was really peculiar, or whether we were just feeling the anxiety that occurs toward the end of every sunspot cycle as solar physicists await the first new-cycle active regions.  We are all familiar with this anxiety.  Flare researchers become anxious because they have contracts to study solar flares.  Energetic particle researchers become anxious because they want more data.  NASA managers become anxious because their spacecraft missions were justified in terms of what they would learn about solar activity.  But most anxious of all are the scientists who predicted the strength of the next sunspot cycle.  So when you encounter the forecasters, appreciate the stress they are under and be kind.  The level of anxiety increased during the 1976 minimum when Jack Eddy reminded us that sunspots became particularly scarce during the 70-year interval 1645-1715 and that another interruption of the sunspot cycle might occur at any time \\citep{Eddy_76}.  His talks alarmed some people, and sent reporters to solar observatories to find out if we were headed into another Maunder Minimum.  We did not enter another Maunder Minimum in 1976 and we have emerged unscathed from two subsequent minima since that time.  Let's see how the present sunspot minimum compares with some of these earlier ones. ",
        "conclusions": ""
    },
    "1005/1005.1831_arXiv.txt": {
        "abstract": "We present a new model for the spectral evolution of Pulsar Wind Nebulae inside Supernova Remnants. The model couples the long-term dynamics of these systems, as derived in the 1-D approximation, with a 1-zone description of the spectral evolution of the emitting plasma. Our goal is to provide a simplified theoretical description that can be used as a tool to put constraints on unknown properties of PWN-SNR systems: a piece of work that is preliminary to any more accurate and sophisticated modeling. In the present paper we apply the newly developed model to a few objects of different ages and luminosities. We find that an injection spectrum in the form of a broken-power law gives a satisfactory description of the emission for all the systems we consider. More surprisingly, we also find that the intrinsic spectral break turns out to be at a similar energy for all sources, in spite of the differences mentioned above. We discuss the implications of our findings on the workings of pulsar magnetospheres, pair multiplicity and on the particle acceleration mechanism(s) that might be at work at the pulsar wind termination shock. ",
        "introduction": "\\label{sec:int} Pulsars (PSRs) are known to release most of their rotational energy in the form of a relativistic magnetized wind, whose particle component is mostly made of electron-positron pairs, with, possibly, a minority of ions, while the magnetic field is almost purely toroidal far away from the light cylinder (e.g. \\cite{goldreich69}). The wind is initially confined by the slowly expanding Supernova Remnant (SNR), which leads to the formation of a termination shock at some distance from the pulsar, and of a bubble of relativistically hot fluid beyond it. This is the so-called Pulsar Wind Nebula (PWN), that shines through synchrotron and Inverse-Compton emission from radio wavelengths up to $\\gamma$-rays.  The typical energy that a pulsar injects into the PWN during its entire lifetime is about $10^{-2}$ of the typical energy of the SN explosion ($10^{51}$ erg). Therefore, the presence of an energetic PSR, has little effect on the global evolution of the SNR. In contrast, the evolution of the PWN strongly depends on its interaction with the SNR (see Sec.~\\ref{sec:neb}).  Modeling of the evolution and dynamical properties of PWN-SNR systems has much improved since the pioneering one-dimensional analytical works by \\citet{rees74} and \\citet{kennel84a}. In recent years, the impressive advances in X-ray observations, due to {\\it CHANDRA} and {\\it XMM-Newton}, have led to a renewed interest in these objects. A major effort has been devoted to explaining the observed X-ray properties of the Crab Nebula and other young systems through multi-dimensional time-dependent approaches to their dynamics \\citep{komiss04,ldz04,ldz05}. Several authors \\citep{blo01,swa01,me04b,dejager08} have presented numerical evolutionary models, extended to advanced ages (typically up to $\\sim 10^5$ yr), and to the case of fast moving pulsars \\citep{swa04}, but in fact, fully multidimensional models are still lacking in these cases. The reason for this is that detailed modeling of older systems is far more difficult than for the young ones: the latter are brighter, hence providing us with high quality multi-frequency data, and are often better constrained in terms of the most relevant physical quantities, such as age, distance and pulsar energetics.  The overall evolution of old objects is much more complex than the smooth expansion appropriate to describe young systems. In general, even within the most simplistic approach, the long-term evolution of a PWN-SNR system depends on many parameters: the SN energy, the mass of the ejecta, their density structure, the density of the ISM, the PSR luminosity, its spin-down history \\citep{me04b}. A model for particle injection and spectral evolution must be added to all of this, if the aim is that of deriving information from direct comparison with observations. At present, multi-dimensional numerical simulations are not suitable to properly investigate the large parameter space that older systems may span.  The situation is actually delicate also for many young systems. Indeed, accurate multidimensional results, including emission maps to be directly compared with observations, have been so far presented only for the Crab Nebula. But this is a very special object, for which measurements are available through almost the entire electromagnetic spectrum and almost all the aforementioned parameters are known, including the age and the pulsar rotational history (initial spin frequency and braking index). For other young systems, like 3C58 or MSH 15-52, the quantity and quality of information available is much poorer than in the case of Crab and simplified evolutionary models are again the only possible tool for investigation, at least for a start.  Models for the evolution of the emission properties of PWNe have a long history, from the work by \\citet{pacini73}, to the most recent developments by \\citet{dejager08}, and \\citet{gelfand09}. These later works have attempted at taking into account both the dynamics of PWN-SNR systems, as derived from accurate numerical simulations, and the spectral evolution of the particle distribution function. In particular, the work by \\citet{gelfand09} is exemplary for showing the richness of behaviors that is found even within the one-zone treatment, and the degree of complexity that one can investigate using simple tools.  While constraining the unknown parameters is preliminary to detailed multi-dimensional modeling of the systems, which can only sample a much more limited range of parameters, the question arises of how strong and reliable the results obtained with a 1-zone approach are. The recent success of 2-d models of PWNe at reproducing so many of the observed features has interesting consequences in this respect, in that it lends support to the 1-zone description over the 1-dimensional approximation. The latter  inevitably leads to onion-like structures \\citep{kennel84b}, where particles injected at earlier times are located at larger radii, with a one-to-one correspondence between the properties of the particle distribution function and the distance from the central PSR. Comparison of multidimensional models with observations has shown that such ordered structure is not what is realized in PWNe, where an overall turbulent flow is injected at the termination shock and survives through most of the nebular extent \\citep{camu09}. For the typical flow times, that are much smaller than the PWN age, this turbulent flow-structure leads to efficient mixing, which is what a 1-zone description is based on. This also finds some direct hints from observations where detailed data are available, as in the case of the Crab Nebula: the remarkable homogeneity of radio spectral index across the nebula, and the extent of the X-ray emitting region which exceeds the predictions of any simple 1-dimensional flow model in the absence of particle diffusion far more efficient than the Bohm rate \\citep{amato00}.  In the following we present a simplified evolutionary model, similar in spirit to the work by  \\citet{gelfand09}. Our approach combines analytical results and numerical simulations for the description of the dynamics of the PWN-SNR system and the evolution of the particle distribution function within the PWN: the key features of the interaction between a PWN and the surrounding SNR are adequately reproduced during different evolutionary stages. The model is here applied to a few systems of various ages, both young and older ones with the aim of showing how this kind of modeling can help the extraction of information on the system properties from multi-wavelength observations: we try to clarify how these properties depend on the different parameters, also discussing potential degeneracies and the observations that would best serve to disentangle them.  Special attention is devoted to the particularly interesting and difficult task of constraining the particle injection in PWNe. The relevance of this issue in terms of pulsar physics and shock acceleration physics deserves a somewhat extended discussion, which will be the topic of next section (Sec.~\\ref{sec:mul}).  The rest of the paper is organized as follows. Sec.~\\ref{sec:neb}, Sec.~\\ref{sec:par} and Sec.~\\ref{sec:emis} describe, respectively, how the nebular evolution is implemented in our model, the scheme we adopt to follow the evolution of the particle distribution function and the emission model. In Sec.~\\ref{sec:res} we show the application of our model to several objects, and, finally, in Sec.~\\ref{sec:dis} we discuss the implications of our findings. ",
        "conclusions": "\\label{sec:dis} Let us briefly review here our findings, and what general conclusions can be drawn from the results of our attempt at modeling different objects, both young and old.  \\subsection{Summary of the results}  First of all, it is remarkable that our one-zone model, despite its simplifications, has proven able to account for the observed Spectral Energy Distributions in all cases. This is already an interesting result, since the assumption of efficient mixing that is at its base is likely a poor approximation at high energies (in the X-ray band), where the flow pattern and the magnetic field structure in the inner nebula are expected to play a dominant role.  The agreement between the model fit and the data is extremely good in the case of Crab, where the unknown parameters are reduced to a minimum (essentially the SNR properties). Our model seems to work rather well for young systems in general, but the quality of the data and the limited spectral coverage lead to larger uncertainties. In the case of strongly magnetized sources, like 3C58 and the nebula associated to PSR B1509, we find that in order to reproduce the data our model requires an energy input in the PWN which exceeds the pulsar release by 20-30\\%, judging from the current spin-down power. This conclusion does not seem to be affected by uncertainties in the distance to the sources: moving them closer to (further from) us would lower (increase) the estimated brightness, but at the same time it would make them smaller and younger (larger and older), leaving the discrepancy between the radiated and accumulated energy almost unaltered. We think it likely that this discrepancy is due to the simplifying assumptions of the one-zone model, which become progressively more important at large magnetizations. In particular, as also shown by recent multidimensional simulations \\citep{vol08}, the radio and X-ray emitting particles might sample different magnetic fields. At the same time it is worth keeping in mind that the estimate of the PSR spin-down power is based on the assumed canonical moment of inertia for neutron stars: $I=10^{45}$ g cm$^2$. A 20\\% discrepancy in two systems is compatible with present uncertainties.  As one can easily see from Tab.~\\ref{tab:1}, given the already high value of the quantity $\\eta_e+\\eta_M$, there is little energy ($\\sim 20$\\% at most) left to be carried by any other higher energy component (ions or leptons in the return current). The idea that a particle population with a larger Larmour radius could be carrying most of the wind energy was made attractive  by the following three main reasons: these particles could help solve the problem of particle acceleration at the pulsar wind termination shock; they could explain the wisps variability in the Crab Nebula; finally, in the case of ions, they could solve the discrepancy between the predictions of 1-d MHD models and the gamma-ray flux observed from Crab.  As far as particle acceleration is concerned, at a relativistic pair shock, efficient acceleration is shown to be possible only if the shock magnetization is extremely low \\citep{spitkovsky08}. The 2-D MHD models of PWNe \\citep{ldz05,camu09}, use a simple model of field reversal across the equator where a current sheet occurs, such that the appropriate conditions for acceleration are thought to be realized only in a small sector of the pulsar wind termination shock, with no more than a few percent of the total wind energy flux flowing through it.  On the other hand, the presence of an energetically significant ion component in the pulsar wind is found to lead to efficient acceleration of the pairs \\citep{hoshino92,amato06}. This effect is the consequence of the large Larmor radius of the ions, whose gyration introduces long wavelength turbulence through cyclotron instability. The same result could come from the presence of high energy leptons, accelerated as a consequence of runaway dynamics in the equatorial current layer (Arons, in preparation).  As to the wisp variability, \\citet{spitkovsky04} showed that this could be explained as the result of compression waves associated with the high energy current carriers' (ions, in their case) gyration in the shock region: in order to reproduce the observed brightness contrast, again one would need the current carriers (the ``beam'') to be energetically significant: $\\gamma_{beam} m_{beam}/\\gamma_w m_\\pm > 2 \\kappa_\\pm$. An alternate explanation of the wisp variability, which does without the kinetic effects of the high energy particles beam, has been recently shown by \\citet{camu09}, who have proved that it can be recovered also within the framework of pure MHD, due to global instabilities of the termination shock.  One final consideration concerns the possible detection of signatures of high energy protons. Modeling of the TeV emission from the Crab Nebula, based on the 1D Kennel \\& Coroniti flow \\citep{atoyan96} was shown to underpredict the observed flux by a factor $\\sim 5$, leading to suggest a possible contribution from the decay of neutral pions produced in nuclear collisions of relativistic protons. More recently, the presence of relativistic protons as the source of high-energy gamma-ray emission was suggested also in the Vela PWN \\citep{horns06}. In the case of Vela, however, a later determination of the density of the thermal protons (that would serve as a target for nuclear collisions) resulted in too low a value and led to strongly question the initial claim \\citep{lamassa08}. At the same time, \\citet{vol08} showed that using a multidimensional model for the flow structure in Crab, the estimated IC-TeV emission is easily overproduced, for the same parameters that allow one to better fit the lower energy synchrotron emission. This suggests that the results are indeed strongly model-dependent and that the discrepancies are likely due to problems with the adopted MHD model.  This is also the conclusion we reach in this work. We can reproduce the TeV emission from Crab, and find results in agreement with {\\it Fermi} observations without any need for a proton contribution. In the case of both B1509 and 3C58, our pair+magnetic field energetic is already exceeding by 20\\% the estimated PSR energy input, and fitting the gamma-rays through $\\pi^0$ decay requires that protons alone carry far more energy than the nebula is currently estimated to store. Finally, in the case of Kes75, in order to ascribe the excess TeV emission to protons, one needs to assume that protons can effectively diffuse out of the nebula and experience the whole ejecta mass as a target.  Our model seems to constrain with good accuracy the value of $\\epsilon_c$. It is not possible to vary the value of $\\epsilon_c$ by more than a factor of a few still satisfying the overall energetics and reproducing the radio and X-ray data. At the same time, our model allows us to derive, for each object, an upper limit on the wind Lorentz factor and a lower limit on the pair multiplicity. The limits we put on $\\gamma_w$ and $\\kappa$ are strictly valid only if the wind has a unique Lorentz factor and the energy scales $\\epsilon_m$ and $\\epsilon_c$ respect the assumed scaling with time, proportional to the pulsar voltage. It is not easy to predict how a possible latitude dependence of the wind Lorentz factor, such as that included in axisymmetric models of PWNe ({\\it e.g.} \\citet{ldz04}), or a different dependence on time of the energy scales would affect our conclusions.    \\subsection{Interpretation of spectral breaks}  Our study has reinforced a long standing puzzle: the electron (and positron) accelerator at work in PWNe knows how to create a spectrum convex in energy space, best represented by a broken power law.  As was advanced by \\citet{kennel84b}, the spectral steepening observed between optical/soft X-ray energies and the harder ($\\varepsilon > 10$ keV) spectrum is well understood as the effect of synchrotron energy losses, with the pre-cooled spectrum a power law $N(E) \\propto E^{-\\gamma_2}, \\; 2.1 < \\gamma_2 < 2.8$ from Table~\\ref{parameters}. The steepening between mid-infrared and optical, when observed, requires energy space structure in the accelerator (an ``intrinsic break''). The high energy spectrum can be attributed to diffusive acceleration at the termination shock, at least qualitatively - Fermi acceleration in relativistic very low sigma shocks in the test particle limit yields a spectrum $N(E) \\propto E^{-2.2}$ \\citep{keshet05}. The radio spectrum requires a much harder distribution of the lower energy particles, $N(E) \\propto E^{-\\gamma_1}, \\; 1.2 < \\gamma_1 <1.7$.  The standard model \\citet{kennel84a, kennel84b}, developed to account for the radiation from Crab at near infrared and shorter wavelengths, assigns the conversion of the pulsar wind energy in the particle spectra that emit the observed synchrotron radiation to diffusive Fermi acceleration at the wind termination shock. Diffusive shock acceleration always shows particle spectra with a Maxwellian at low energy, plus a power law supra-thermal tail at high energy, with the temperature of the Maxwellian set by the shock jump conditions (except when the acceleration of the tail is very efficient). These properties are well exhibited by Particle-in-Cell simulations of relativistic shocks in {\\it unmagnetized} $e^\\pm$ plasmas (or, if magnetized, in upstream quasi-parallel flow geometry \\citep{spitkovsky08,  sironi09}).  Our models identify the transition between the soft, high energy spectrum and the hard, low energy spectrum as being at the energy $\\epsilon_c$, typically $10^5 - 10^6m_\\pm c^2$ (see  Table~\\ref{parameters}). In the \\citep{kennel84b} model, this energy was identified as the lower cutoff of the shock accelerated power law, and that energy was identified as the ``temperature'' $\\approx \\gamma_w m_\\pm c^2$ of the flow downstream of the termination shock, thus giving rise to the belief that pulsar winds have upstream flow Lorentz factor $\\gamma_w \\sim 10^6$.  The Kennel and Coroniti model and its descendants deliberately neglected the radio emission from the Crab Nebula (and, by extension, other PWNe.) The large (by number) population of particles with $E < \\epsilon_c$ suggests $\\gamma_w$ to be much smaller, if the radio and mid- to far-IR emitting particles come from the pulsar. That the pulsars are the most likely source of the low energy particles in each of the nebulae gains support from the existence of the radio ``wisp'' features in the Crab \\citep{bie04} closely associated with the similar time variable structures seen in optical and X-ray imaging, as well as similar structure seen in 3C58 \\citep{bie06} - those structures are clearly coincident with the termination shock. Then, in the regions of the flow populated by the large particle flux feeding the low energy population, $\\gamma_w \\sim 10^4$ or smaller.  The low energy particle spectrum definitely is not a single temperature relativistic Maxwellian, with $T \\sim \\epsilon_c$.  Since the termination shock is not spherical, becoming more oblique in higher rotational latitude regions \\citep{komiss04, ldz04}, the post-shock temperature declines with increasing latitude, suggesting the low energy particles enter in the polar regions of the outflow. However, replicating the observed spectrum as the envelope of a sequence of Maxwellians requires a factor of $\\sim 100$ variation in temperature, which requires most of the mass flux being nearest the rotation poles. There is no sign of pole to equator asymmetry in the radio emission near the pulsars, except for the radio wisp features in the Crab, which are components of the immediate post-shock flow, as judged from the optical and harder photon emission. However we cannot rule out that mixing and effects related to integration along the line of sight might prevent the detection of the implied variations.  However, let us assume in the following, that also the low-energy particle spectrum results from some acceleration process, rather than from the convolution of different thermal distributions, and let us speculate on the nature of such a process.  Turbulence and associated Fermi II acceleration in and around the termination shock is one appealing possibility. The \"wisp\" motion, observed from the radio through X-rays (see \\citet{hes08} for a comprehensive review) has recently been interpreted as the result of the strongly variable termination shock structure found in high resolution MHD simulations of the Crab Nebula \\citep{camu09}. The shock instability implied is a termination shock variant of the Standing Accretion Shock Instability, with outer scale variable velocity $\\delta v \\sim v \\sim 0.6c$ and length scale $\\delta r \\sim r\\sim 0.5-1$ light years. The magnetized motions observed in the simulations (which do well in replicating the time variable spatial structure observed in the nebular ``wisps'') can act as an accelerator through the Fermi II process \\citep{kar62, staw08}, creating electron and positron spectra $N(E) = N_0 (E/E_0)^{-s}$.  In quasi-linear models of Fermi II acceleration in isotropic small amplitude Alfven turbulence ({\\it e.g.} \\citet{staw08}), $s= \\sqrt{9/4 + \\epsilon} -1/2$, where $\\epsilon = T_{accel}/T_{escape}$. These models include an analogue of scattering from large amplitude magnetic inhomogeneities (Fermi's original model), in the case where the wave energy spectrum scales in proportion to $k^{-2}$, with $k$ the wavenumber. This case is germane to acceleration due to pair interaction with large scale moving magnetic fluctuations: the acceleration time scale for interaction with large scale moving ``eddies'' is $T_{accel} \\approx \\lambda_0/c (c/v_{eddy})^2 (B/\\delta B)^2 \\sim$, where $\\lambda_0$ is the outer scale of the turbulence, comparable to the shock radius in our case \\citep{camu09}. The high downstream flow velocity suggests particles escape the turbulence zone through advective loss rather than diffusive escape (at energies much less than a PeV, microscopic diffusion across B, even at the Bohm rate, is negligible compared to the advective losses). Then at flow speed $\\sim c/3$, $T_{escape} \\sim 3 R_{shock} /c$, leading to $\\epsilon \\sim (1/3) (B/ \\delta B )^2 (\\lambda_0 /R_{shock})$. Taking the parameters to be unity yields $\\epsilon \\sim 1/3$, in which case the accelerated spectrum is very hard: $N(E) \\propto E^{-1.3}$ which is similar to the spectrum inferred for the radio emitting pairs in PWNe.  Because of the strong radial mixing observed in the MHD models of PWNe, the spectrum created in this turbulent acceleration zone will fill the whole body of the synchrotron emitting nebulae, which provides a natural explanation of the observed lack of gradients in the radio spectral index of the Crab \\citep{bie92}. Thus, we revive, in a modern form, the suggestion of \\citet{barnes73}, that the wisp motions are responsible for a part of the particle acceleration required to account for PWNe synchrotron emission.  Turning this idea into a full physical model may require extending the MHD models to three dimensions. All models to date have been 2D, axisymmetric, with exclusively toroidal magnetic field.  Particles then cannot scatter radially, in order to sample adjacent moving eddies, but are confined to toroidal flux tubes, which move in and out in radius around average positions in the outflow (and inflow, at higher latitudes) in the turbulent region.  It is possible that magnetic pumping in these flux tubes might substitute for particle scattering in quasi-isotropic turbulence as an acceleration mechanism ({\\it e.g.}, \\citet{mel69, kuij97}). However, polarization studies already demonstrate that there are substantial poloidal magnetic fields in PWNe inner regions, at least in the Crab.  The termination shock instability identified by \\citep{camu09} is not likely to be restricted to 2D - perhaps the assumption of quasi-isotropic turbulence, with forcing scale $\\sim R_{shock}$, is the most natural starting point for further pursuit of this idea.  It is clear that a key point for the model to succeed is that the proposed turbulently accelerated spectrum at low energy merges smoothly into the shock accelerated spectrum at high energy. Since the macroscopic turbulence is closely associated with the shock, a smooth merger is at least thinkable. In this context the meaning of $\\epsilon_c$ changes completely with respect to the Kennel and Coroniti interpretation: this should not be identified with the post-shock temperature (which according to the present model is about an order of magnitude lower), but rather with the energy above which the low energy accelerator (Fermi II in the turbulent flow, in the suggestion made here) fades out and relativistic DSA takes over.  \\subsection{How Do We Move Forward?}  Multi-D high resolution MHD and PIC simulations of the termination shock region separating the freely expanding wind from the subsonically expanding nebulae will shed light on the possibly turbulent flow, and on whether that turbulence can act as the low energy accelerator discussed previously. A useful first step would be to couple the diffusion-advection equation in energy space to the MHD calculations. Observationally, infrared observations of the young PWNe, using the Herschel space telescope \\citep{pilbratt09} as well as near infrared instruments on the ground and in space, will greatly improve our understanding of the physics behind the broken power laws in energy space inferred here for the injected particle distributions. These instruments will have more than sufficient angular resolution to resolve the PWNe, thus emphasizing the need for multidimensional models to quantitatively interpret the observations.  For all the objects studied, the plasma injection rates, here inferred from the evolutionary models, not simply taken from the average over the nebular lifetimes, exceed $10^5$ times the Goldreich-Julian rate. Such high rates can be understood only as the result of pair creation within the pulsars' magnetospheres - given the spectral continuity in the SEDs, models which rely on the radio and far-IR emitting particles, which dominate the injection rates, being accelerated from the thermal plasma in and around the nebulae are less likely than injection from the pulsars. Our results support and extend a long held suspicion \\citep{gallant02}, previously based only on analysis of the Crab Nebula, that the particle loss rates from young, high voltage pulsars are substantially in excess of the inferences of particle {\\it outflow} from all known magnetospheric pair creation models. The answer may lie in intrinsic time dependence of pair creation within pulsar polar caps ({\\it e.g.} \\citet{timok09}), or in the so far unexplored outer magnetosphere accelerators associated with the boundary layer return currents separating the closed and open field lines \\citep{aro09}. {\\it Fermi} gamma ray observations are and will be useful in constraining magnetospheric pair creation models."
    },
    "1005/1005.4249_arXiv.txt": {
        "abstract": "We apply the Constitution compilation of 397 supernova Ia, the baryon acoustic oscillation measurements including the $A$ parameter, the distance ratio and the radial data, the five-year Wilkinson microwave anisotropy probe and the Hubble parameter data to study the geometry of the universe and the property of dark energy by using the popular Chevallier-Polarski-Linder and Jassal-Bagla-Padmanabhan parameterizations. We compare the simple $\\chi^2$ method of joined contour estimation and the Monte Carlo Markov chain method, and find that it is necessary to make the marginalized analysis on the error estimation. The probabilities of $\\Omega_k$ and $w_a$ in the Chevallier-Polarski-Linder model are skew distributions, and the marginalized $1\\sigma$ errors are $\\Omega_m=0.279^{+0.015}_{-0.008}$, $\\Omega_k=0.005^{+0.006}_{-0.011}$, $w_0=-1.05^{+0.23}_{-0.06}$, and $w_a=0.5^{+0.3}_{-1.5}$. For the Jassal-Bagla-Padmanabhan model, the marginalized $1\\sigma$ errors are $\\Omega_m=0.281^{+0.015}_{-0.01}$, $\\Omega_k=0.000^{+0.007}_{-0.006}$, $w_0=-0.96^{+0.25}_{-0.18}$, and $w_a=-0.6^{+1.9}_{-1.6}$. The equation of state parameter $w(z)$ of dark energy is negative in the redshift range $0\\le z\\le 2$ at more than $3\\sigma$ level. The flat $\\Lambda$CDM model is consistent with the current observational data at the $1\\sigma$ level. ",
        "introduction": "The accelerating expansion of the universe was first discovered by the type Ia supernova (SN Ia) observations \\cite{acc1,acc2}. The phenomena of acceleration could be explained straightforwardly by introducing an exotic source of matter with negative pressure, the so-called dark energy, which dominates the total matter content of the universe at the present epoch and causes the expansion to accelerate. During the past decade, in addition to the simple cosmological constant model, a lot of dynamical dark energy models, such as the quintessence \\cite{quint}, phantom \\cite{phantom}, k-essence \\cite{k}, tachyon \\cite{tachyonic}, quintom \\cite{Feng:2004ad}, h-essence \\cite{hessence}, Chaplygin gas \\cite{chaplygin}, holographic dark energy \\cite{holo}, $f(R)$ \\cite{fr}, Dvali-Gabadadze-Porrati \\cite{dgp} models, etc, have been proposed. Although a lot of efforts have been made to understand the driving force of the accelerating expansion and the property of dark energy, whether dark energy is dynamical or not is still an open question. Therefore, it is necessary to study the nature of dark energy such as the evolutions of its energy density and equation of state.  Apart from phenomenological models, another effective approach to study dark energy is through the observational data. Recently, based on the popular Chevallier-Polarski-Linder (CPL) parametrization of dark energy  \\cite{cpl}, it was found that the flat $\\Lambda$CDM model is inconsistent with the current data at more than $1\\sigma$ level \\cite{star,huang}. In \\cite{star}, it was suggested that the cosmic acceleration is slowing down from $z\\sim 0.3$. In \\cite{huang}, it was claimed that dark energy suddenly emerged at redshift $z\\sim 0.3$. Furthermore, possible oscillating behavior of dark energy was found in \\cite{cai}. However, no evidence for dark energy dynamics was found  in \\cite{corray9,gong,gong10}. It was argued that the systematics in different data sets heavily affected the fitting results from observational data \\cite{gong,gong10}. To further study the dynamics of dark energy, it is necessary to apply more complimentary observational data. In this paper, we combine the Constitution sample of 397 SN Ia data \\cite{consta}, the model independent $A$ parameter from the baryon acoustic oscillation (BAO) measurements \\cite{bar}, the two BAO distance ratios at $z=0.2$ and $z=0.35$ \\cite{wjp}, the radial BAO measurements at $z=0.24$ and $z=0.43$ \\cite{eg}, the five-year Wilkinson microwave anisotropy probe data (WMAP5) \\cite{kdn}, and the Hubble parameter $H(z)$ data \\cite{hz1,hz2} to probe the geometry of the universe and the nature of dark energy by using the CPL and Jassal-Bagla-Padmanabhan (JBP) \\cite{jbp} parameterizations. We first use the simple $\\chi^2$ method of joined contour estimation to obtain the constraints on the model parameters. However, the simple $\\chi^2$ method by fixing other parameters at their best fit values has some drawbacks because we neglect the correlation effects between the parameters and the degeneracy between parameters was not considered. When the parameters are strongly correlated, the errors of some parameters will be under-estimated if we fix the other parameters at their best fit values. So we also apply the Monte Carlo Markov chain (MCMC) method to obtain the marginalized errors of the model parameters. The advantage of the MCMC method is that it considers the correlations between the model parameters and the result is more reliable.  The paper is organized as follows. In section 2, we present the SN Ia data \\cite{consta}, the BAO data \\cite{bar,wjp,eg}, the WMAP5 data \\cite{kdn} and the $H(z)$ data, and all the formulas related with these data. In section 3, We use the $\\Lambda$CDM model as an example to show how to apply the data to constrain cosmological models. In section 4, we use the CPL model to study the geometry of the universe and the property of dark energy. The JBP model is used to probe the geometry of the universe and the evolution of dark energy in section 5. We conclude the paper in section 6. ",
        "conclusions": "Applying the simple $\\chi^2$ method, we fitted the CPL and JBP models to the combined SN Ia, BAO, WMAP5 and $H(z)$ data, and obtained the constraint on the property of dark energy. In both CPL and JBP models, there are four parameters $\\mathbf{p}=(\\Omega_{m},\\ \\Omega_{k},  \\ w_0, \\ w_a)$. When we apply the BAO and WMAP5 data, we need to add two more parameters $\\Omega_b h^2$ and $h$. If we make the joint error analysis, we have six parameters and it will be hard to get a good joint constraint on all these parameters. Therefore, we take $\\Omega_b h^2=0.02273$, and find out the best fit values of the parameters $\\mathbf{p}$ and $h$ which minimize $\\chi^2$; then we fix the parameters $\\Omega_m$, $\\Omega_k$ and $h$ at their best fit values to obtain the joint constraints on $w_0$ and $w_a$. For the CPL model, the contours of $w_0$ and $w_a$ (see figure \\ref{cpl1:a}) show that the $\\Lambda$CDM model is excluded at more than $3\\sigma$ level. The JBP model is consistent with the $\\Lambda$CDM model at the $1\\sigma$ level. Since we get the contours of $w_0$ and $w_a$ by fixing the other parameters at their best fit values, we neglect the correlation effects of the parameters and the conclusion based on this method may not be reliable. To confirm this, we use the MCMC method to analyze the CPL and JBP models and obtain the marginalized probabilities of the parameters. For the CPL model, the probability distributions of $\\Omega_k$ and $w_a$ are skew distributions, and the marginalized $1\\sigma$ errors are $\\Omega_m=0.279^{+0.015}_{-0.008}$, $\\Omega_k=0.005^{+0.006}_{-0.011}$, $w_0=-1.05^{+0.23}_{-0.06}$ and $w_a=0.5^{+0.3}_{-1.5}$. In the CPL model, the probability distributions of $\\Omega_k$ has a local maximum in addition to a global maximum. The uncertainties in $\\Omega_k$ and the degeneracies between $\\Omega_k$, $w_0$ and $w_a$ lead to under-estimation of the error contours of $w_0$ and $w_a$ if we fix $\\Omega_k$ at its global best fit value, and the wrong conclusion that the $\\Lambda$CDM model is excluded at more than $3\\sigma$ level. However, this does not happen for the JBP model. For the JBP model, the parameters have Gaussian distributions, and the marginalized $1\\sigma$ errors are $\\Omega_m=0.281^{+0.015}_{-0.01}$, $\\Omega_k=0.000^{+0.007}_{-0.006}$, $w_0=-0.96^{+0.25}_{-0.18}$ and $w_a=-0.6^{+1.9}_{-1.6}$.  In summary, in addition to use the usual SN Ia, BAO $A$ or BAO distance ratio, and WMAP data, we also use the radial BAO measurements and the $H(z)$ data to fit the CPL and JBP models. We find that the equation of state parameter of dark energy $w(z)<0$ at more than $3\\sigma$ level in the redshift range $0\\le z\\le 2$, and the flat $\\Lambda$CDM model is consistent with the current observational data at the $1\\sigma$ level. Furthermore, we find that we need to do the marginalized analysis to estimate the errors of the model parameters.   \\ack  NP was partially supported by the project A2008-58 of the Scientific Research Foundation of Chongqing University of Posts and Telecommunications, and the NNSF of China under grant no 10947178. YG was partially supported by the NNSF key project of China under grant no 10935013, the National Basic Research Program of China under grant no 2010CB833004, and the Natural Science Foundation Project of CQ CSTC under grant no 2009BA4050. ZZ was partially supported by the NNSF Distinguished Young Scholar project under Grant no 10825313, and the National Basic Research Program of China under grant no 2007CB815401."
    },
    "1005/1005.4763_arXiv.txt": {
        "abstract": "Timing analysis of the {\\it INTEGRAL}-IBIS and \\emph{Swift}-BAT light curves of the Supergiant Fast X-ray Transient (SFXT) IGR J16465--4507 has identified a period of 30.32$\\pm$0.02\\,days which we interpret as the orbital period of the binary system.  In addition 11 outbursts (9 of which are previously unpublished) have been found between MJD 52652 to MJD 54764, all of which occur close to the region of the orbit we regard as periastron. From the reported flux outbursts, we found a dynamical range in the interval $\\sim$30-80. Although in this regard IGR J16465--4507 cannot be considered a  classical SFXT for which typical dinamical ranges are $>$100, still our reported values are significantly greater than  that of classical persistent variable supergiant HMXBs ($<20$), supporting  the idea that IGRJ16465--4507 is an intermediate SFXT system, much like few other similar cases reported in the literature. ",
        "introduction": "During the last few years, the IBIS instrument \\citep{Ubertini:2003kn} on board the \\emph{INTEGRAL} gamma-ray satellite \\citep{Winkler:2003hh} has played a key role in discovering a previously unrecognized class of high mass X-ray binaries which display a very unusual fast X-ray transient behaviour and host a massive OB supergiant star as companion: the so-called Supergiant Fast X-ray Transients (SFXTs; \\citealt{Zand:2005dp,Sguera:2005os,Sguera:2006ns,Negueruela:2006lt}). To date approximately 10 show clear characteristics of being a SFXT (plus the same number of candidates), characterized by long periods of low X-ray activity level (with luminosities or upper limits typically in the range 10$^{32}$-10$^{34}$ erg s$^{-1}$), interspersed with short, strong, flares lasting from a few hours to no more than a few days and reaching peak luminosities of 10$^{36}$-10$^{37}$ erg s$^{-1}$. The identification of periodicity in the SFXT outburst behaviour, most likely representing the orbital period, is a key diagnostic for studying the geometry of these systems and hence for understanding the still unknown physical mechanisms behind their very unusual behaviour (see \\citealt{2009AdSpR..43.1464S} for a review). Because the fast outbursts from SFXTs are relatively rare occurrences, the discovery of such periodicities is a very challenging task which has been fulfilled only for an handful of SFXTs, mainly  through long-term temporal monitoring by \\emph{INTEGRAL} (e.g. \\citealt{2009MNRAS.399L.113C,Bird:2009yq, 2009MNRAS.397L..11J, 2009A&A...493L...1Z, 2007A&A...467..249S, 2007A&A...476.1307S}).  IGR J16465--4507 classification as an SFXT is debatable. Its initial discovery on 6 September 2004 was triggered by the \\emph{INTEGRAL} observation in the Norma region of a faint $\\sim$9 mCrab source (18-60 keV) which produced a strong $\\sim$28 mCrab flare on 7 September 2004 before returning to its undetectable pre-flare level \\citep{2004ATel..329....1L}, demonstrating its fast X-ray transient nature. However since then, no more outburst activity has been reported from the source and its temporal and spectral behaviour above 20\\,keV is largely unknown. A \\emph{XMM-Newton} observation performed a few days after the source discovery allowed the detection of high intrinsic absorption ($\\sim$6$\\times$10$^{23}$ cm$^{-2}$) as well as a $\\sim$228\\,s period which was interpreted as the spin period of  a neutron star \\citep{Walter:2006fv,2005A&A...444..821L}. A single bright infrared object falls in the refined \\emph{XMM-Newton} position \\citep{2004ATel..336....1H}, optical and infrared spectroscopy classified it as a supergiant star located at a distance of $9.5^{+14.1}_{-5.7}$\\,kpc, although different spectral types have been proposed (O9.5Ia by \\citealt{2008A&A...486..911N} or B0.5I by  \\citealt{Negueruela:2006lt, Rahoui:2008ig}). \\citet{Walter:2007ww} reported a 22--50 keV quiescent flux of $7.7\\times10^{-12}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ and suggest that IGR J16465--4507 is not a classical SFXT (for which the dynamical range is $>$100, \\citealt{2010arXiv1005.1995C}) but a kind of intermediate SFXT still having a dynamic range greater than that of a classical persistent variable supergiant HMXB (for which the dynamic range is $<$20) such as Vela X-1 \\citep{2008A&A...492..511K}. In comparison a very recent paper by \\citet{2010MNRAS.tmpL..63L} find no evidence of flaring activity from IGR J16465--4507 and inferred  a dynamical range $<$10.  Here we report a comprehensive temporal study of IGR J16465--4507, based on long term \\emph{INTEGRAL} monitoring, which led to the discovery of 9 previously unpublished outbursts as well as the $\\sim$30 days periodicity found by \\citet{2010MNRAS.tmpL..63L}, which we interpret as the binary period. ",
        "conclusions": "IGR J16465--4507 has been shown to exhibit an $\\sim$30 day periodicity in both the {\\em INTEGRAL}-IBIS and {\\em Swift}-BAT light curves. This period is attributed to the orbital period of the system and corresponds to a semi-major axis of 121.9-126.3$R_{\\sun}$.  This system is remarkably similar to SAX J1818.6--1703, except that IGR J16465--4507 is known to be further away, helping explain the lack of underlying emission that can only be seen in the phase-folded light curve of the entire {\\em INTEGRAL} data set and not in the recurrence analysis. A low dynamic range is observed, which is higher than that of a typical supergiant HMXB. This would suggest that IGR J16465--4507 is an intermediate SFXT and that these may form a continuum ranging from typical supergiant HMXBs to the extreme SFXTs."
    },
    "1005/1005.4280_arXiv.txt": {
        "abstract": "Dark matter (DM) is currently searched for with a variety of detection strategies. Accelerator searches are particularly promising, but even if Weakly Interacting Massive Particles (WIMPs) are found at the Large Hadron Collider (LHC), it will be difficult to prove that they constitute the bulk of the DM in the Universe $ \\Omega_{\\rm \\tiny{DM}} $. We show that a significantly better reconstruction of the DM properties can be obtained with a combined analysis of LHC and direct detection (DD) data, by making a simple Ansatz on the WIMP local density $ \\rho_{\\neut} $, i.e., by assuming that the local densiy scales with the cosmological relic abundance, $ (\\rho_{\\neut}/\\rho_{\\rm \\tiny{DM}})=(\\relic/\\Omega_{\\rm \\tiny{DM}} $). We demonstrate this method in an explicit example in the context of a 24-parameter supersymmetric model, with a neutralino LSP in the stau co-annihilation region. Our results show that future ton-scale DD experiments will allow to break degeneracies in the SUSY parameter space and achieve a significantly better reconstruction of the neutralino composition and its relic density than with LHC data alone. ",
        "introduction": "\\label{sec:introduction} Identifying the nature of the dark matter (DM) remains one of the central unsolved problems in modern particle physics and cosmology. A generic Weakly Interacting Massive Particle (WIMP) is among the best-motivated possibilities since it can be thermally produced in the early Universe in the right amount to account for the observed DM density. Indeed, many theories for Physics beyond the Standard Model contain viable WIMP candidates, as is the case of Supersymmetry (SUSY) when the lightest SUSY particle (LSP) is the lightest neutralino  (a linear superposition of the supersymmetric partners of the gauge and Higgs bosons) \\cite{Jungman:1995df,Munoz:2003gx,Bertone:2004pz,book}.  DM can be searched for in various ways. One possibility is attempting a {\\it direct} detection, through its scattering off nuclei inside an underground detector. Many experiments have been running or are under construction which are mostly sensitive to the spin-independent part of the WIMP-nucleus cross section, $ \\crosssection $. Among these, the DAMA/LIBRA collaboration reported a possible DM signal \\cite{Bernabei:2008yi,Bottino:2008mf}. However, its interpretation in terms of the elastic scattering of a WIMP with a mass around $ 10-100 \\mbox{ GeV} $ and $ \\crosssection \\sim 10^{-3}-10^{-5}\\, \\mbox{pb} $ has been challenged by other experiments, such as the CoGeNT \\cite{Aalseth:2008rx,Aalseth:2010vx}, CDMS \\cite{Ahmed:2008eu,Ahmed:2009zw}, XENON \\cite{Angle:2007uj} and ZEPLIN \\cite{Lebedenko:2009xe}. The CoGeNT collaboration has itself recently reported an irreducible excess of low-energy events which could also be understood as due to the scattering of a very light WIMP \\cite{Aalseth:2010vx} (see also Ref.\\,\\cite{Chang:2010yk}), but this intepretation has in turn be put under pressure by the XENON-100 results, obtained with a fiducial target mass of 40 kg and 11 days of exposure~\\cite{Aprile:2010um}. Finally, the recent results from the CDMS-II collaboration show two events compatible with a WIMP signal, although these results are still statistically inconclusive \\cite{Ahmed:2009zw}.  The future increase of the sensitivity may clarify the situation, but it is becoming clear that several independent pieces of evidence will be necessary to claim discovery of DM. In fact, even if in principle the WIMP mass and scattering cross section can be determined with some accuracy after its direct detection in one direct detection experiment, provided that the measured event rate is large and the WIMP mass is small \\cite{Green:2007rb,Green:2008rd}, a second direct detection with a different target would actually allow a much more precise determination of the WIMP mass \\cite{Drees:2008bv}, and if the new target is sensitive to the spin-dependent contribution of the WIMP-nucleus cross section it could even be used to discriminate among WIMP candidates \\cite{Bertone:2007xj}.  Another possibility consists in looking for the products of DM annihilation (e.g., high energy neutrinos, gamma-rays or antimatter) and thus {\\it indirectly} reveal the presence of the DM \\cite{Bertone:2004pz,book}. We leave the discussion of this search strategy to a forthcoming work, where we will present the constraints that can be set on the DM parameter space from the observation (or non-observation, see also \\cite{Scott:2009jn}) of DM annihilation radiation \\cite{Bertone:2010}.  Finally, collider experiments, most notably the Large Hadron Collider (LHC), will explore the nature of Physics at the TeV scale, where many of the extensions of the SM that propose DM candidates would manifest themselves. The detection of new Physics in particle colliders can provide crucial information about DM. For example, the mass and spin of the LSP could be determined through the study of kinematic variables \\cite{Cho:2007qv,Cho:2008tj}. However, to prove that the newly discovered particles account for all (or most) of the DM in the Universe, is a challenging task. In fact, although particle accelerators can provide some information about the neutralino relic density \\cite{Baltz:2006fm}, it was found that in many cases the LHC would be unable to determine the precise composition of the neutralino, leading to an unreliable prediction of its relic abundance or to the occurrence of multiple solutions spanning several orders of magnitude, thus not allowing to establish whether or not it is the DM (see also Ref. \\cite{Nath:2010zj} and references therein).  One possibility is to build a new collider, such as the proposed International Linear Collider (ILC), that would allow a much more precise evaluation of the supersymmetric masses and couplings, and a better determination of the inferred relic density, as argued by the authors of Ref.\\,\\cite{Baltz:2006fm}. However, this machine will not be available in the near future, and it is therefore crucial to devise strategies that can be implemented as soon as new particles are discovered at the LHC. Fortunately, direct detection experiments are expected to greatly improve their sensitivity in the next few years and start probing interesting regions of the supersymmetric parameter space. In case of discovery, it will certainly be reassuring if the mass reconstructed from direct detection experiments matched the value obtained from accelerator measurements, since it would prove the existence of a particle which is stable over cosmological timescales. The error on the mass reconstructed from direct detection experiments depends on the DM particle parameters, and on the experimental setup, and the interested reader can find a detailed analysis in Refs.\\cite{Green:2007rb,Green:2008rd}. But one can do much more than checking the compatibility of the two mass determinations. We show here that a combined analysis of the two data sets will allow a much better reconstruction of the DM properties, and a convincing identification of DM particles.  Although the strategy discussed here is model-independent, we work out an explicit example in the context of a 24-parameters supersymmetric model, with a neutralino LSP in the stau co-annihilation region. ",
        "conclusions": "We have investigated the effect of combining information from accelerator searches with data from direct DM detection, assuming realistic measurements at the LHC and in a Germanium detector with an exposure of 300 ton days.  An interesting question is whether the systematic and statistical errors on this quantity and on other relevant physical quantities entering in the calculation of the event rate for direct detection experiments can spoil the reconstruction procedure presented here. For instance, we have assumed a Maxwellian distribution for the velocity dispersion of DM particles, but a more refined analysis should keep into account the uncertainties on this quantity. Fortunately,  recent estimates based on numerical simulations, suggest that the small measured deviations from a Maxwell-Boltzmann distribution lead to errors of 10\\% or less on the recoil rate, and they are therefore subdominant with respect to other uncertainties, such as the error on the nuclear form factor \\cite{Ellis:2008hf}, and especially the error on the observed DM local density. The most important effect of such uncertainties, once marginalized over, would be to widen the pdf's of Fig.~\\ref{fig:omega_vs_sigma} in the vertical direction by less than 10\\%, if one considers only the statistical error on $\\rho_{\\rm \\tiny{DM}}$ derived in Ref. \\cite{Catena:2009mf}, and by up to a factor of two if one also considers the systematic error due to halo triaxiality \\cite{Pato:2010yq}. Since the vertical thickness of the contours in the 2D posterior of Fig.\\,\\ref{fig:omega_vs_sigma} is approximately equal to a factor of 2, we expect that including these uncertainties would not modify qualitatively the marginal posterior distribution for $\\relic h^2$, and our results would still apply. A more detailed discussion of these effects is beyond the scope of this paper, and we leave it to a separate upcoming work. However we explicitly studied the effect of varying the value for the mass of the top, including it as a nuisance parameter in the likelihood. The variation in the reconstructed pdf for the neutralino relic abundance and neutralino-nucleon scattering cross section is negligible.  We stress once more the importance of combining different types of experiments.  The specific case discussed here shows that when reasonable assumptions are made to link the local density to the relic abundance, {\\it a combined analysis of data from accelerators and direct detection experiments allows a significantly better reconstruction of the DM properties}.  This is true in the co-annihilation region discussed here, but it will provide important information for any SUSY scenario, and more in general for any new physics scenario. Even in cases where the LHC data are sufficient to pinpoint the underlying DM scenario, direct detection experiments can corroborate the results, and they can also be used to identify deviations from the standard expansion rate of the Universe at freeze-out that would appear as an inconsistency between the $\\relic $ inferred from LHC data and cosmological measurements.    {\\em Acknowledgements.} D.G.C. is supported by the Ram\\'on y Cajal program of the Spanish MICINN, by the Spanish grants FPA2009-08958, HEPHACOS S2009/ESP-1473 and by the EU network PITN-GA-2009-237920. R.T. would like to thank the EU FP6 Marie Curie Research and Training Network ``UniverseNet'' (MRTN-CT-2006-035863) for partial support, the Instituto de Fisica Teorica (Madrid) and the Institut d'Astrophysique de Paris for hospitality. The work of R. RdA has been supported in part by MEC (Spain) under grant FPA2007-60323, by Generalitat Valenciana under grant PROMETEO/2008/069 and by the Spanish Consolider Ingenio-2010 program PAU (CSD2007-00060). We also thank the support of the spanish MICINN's Consolider-Ingenio 2010 Programme under grant MultiDark CSD2009-00064."
    },
    "1005/1005.1977_arXiv.txt": {
        "abstract": "We report the discovery of the first eclipsing detached double white dwarf (WD) binary. In a pulsation search, the low-mass helium core WD NLTT~11748 was targeted for fast ($\\approx 1$ minute) differential photometry with the Las Cumbres Observatory's Faulkes Telescope North. Rather than pulsations, we discovered $\\approx 180$ s 3\\%--6\\% dips in the photometry. Subsequent radial velocity measurements of the primary white dwarf from the Keck telescope found variations with a semi-amplitude $K_1=271\\pm 3 \\, \\kms$, and confirmed the dips as eclipses caused by an orbiting WD with a mass $M_2=0.648$--$0.771 \\, M_\\odot$ for $M_1=0.1$--$0.2 \\, M_\\odot$. We detect both the primary and secondary eclipses during the $P_{\\rm orb}=5.64$\\,hr orbit and measure the secondary's brightness to be $3.5\\%\\pm 0.3 $\\% of the primary at SDSS-$g^\\prime$. Assuming that the secondary follows the mass--radius relation of a cold C/O WD and including the effects of microlensing in the binary, the primary eclipse yields a primary radius of $R_1=0.043$--$0.039 \\, R_\\odot$ for $M_1=0.1$--$0.2 \\, M_\\odot$, consistent with the theoretically expected values for a helium core WD with a thick, stably burning hydrogen envelope.  Though nearby (at $\\approx150$\\,pc), the gravitational wave strain from NLTT~11748 is likely not adequate for direct detection by the {\\it Laser  Interferometer Space Antenna}. Future observational efforts will determine $M_1$, yielding accurate WD mass--radius measurement of both components, as well as a clearer indication of the binary's fate once contact is reached. ",
        "introduction": "Double white dwarfs (WDs) in tight enough ($P_{\\rm orb}< {\\rm day}$) binaries to reach contact in a Hubble time are expected on theoretical grounds \\citep{nel01}, and are presumed to be the progenitors of highly variable objects: R~CrB stars, AM~CVn binaries, and Type Ia supernovae \\citep{ibe84,web84}. However, examples of these systems are rare, with only 10 known prior to our work \\citep{nel01,bad09,mul09,kil09,mar10,kul10}. Many of these binaries are of immediate interest for the {\\it Laser Interferometer Space Antenna} ({\\it LISA}), providing `verification' sources loud enough to be detected in space \\citep{nel09}.  Spectral measurements of the high proper motion object \\object[NLTT 11748]{NLTT~11748} by \\citet{kaw09} revealed it to be a low-mass ($<0.2 \\, M_\\odot$) helium core WD with $\\log g=6.2 \\pm 0.15$ and $T_{\\rm eff}=8540\\pm 50$\\,K and an H-rich (DA) atmosphere. The strong expectation that low-mass He WDs can only be formed in common envelope events led \\citet{kaw09} to note that this object was likely a binary. Indeed, they found a $40\\, \\kms$ velocity difference between two measured spectra.   NLTT~11748 was originally targeted in a ZZ~Ceti like pulsation search based on theoretical calculations and the $\\log g$ and $T_{\\rm eff}$ measurements \\citep{ste10}. In this failed pulsation search, our rapid  ($\\approx 1$ minute) differential photometry with the Las Cumbres Observatory's Faulkes Telescope North (FTN) revealed $\\approx 180$ s 3\\%--6\\% dips, which were confirmed by Keck spectroscopy to be primary and secondary eclipses from a $\\approx 0.7 \\, M_\\odot$ faint C/O WD companion orbiting at 5.6\\,hr.  NLTT~11748 joins the growing class \\citep{kaw06,kil07,mul09,kil09,mar10,kul10,bad09} of WD binaries where the low-mass WD is a very low-mass ($<0.2 \\, M_\\odot$) helium core WD. These are made when the lower mass star overfills its Roche lobe at the base of the red giant branch \\citep{van96,cal98,bas06a}, triggering a common envelope event that leads to an in-spiral of the now exposed low-mass He core to a shorter orbital period. Theoretical work has found a dichotomy in the evolution of these low-mass He core WDs \\citep{dri99,ser02,pan07} that dramatically impact their cooling rate. For masses $\\gtrsim$0.2\\,$M_{\\odot}$ (metallicity dependent), the H envelope undergoes shell flashes that reduce its mass, eventually allowing the WD to cool rapidly. However, less massive ($<0.2 \\, M_\\odot$) WDs undergo stable H burning for gigayears, dramatically slowing their evolution and keeping them brighter for much longer than expected.  Our observations are summarized in Section \\ref{sec:obs}, demonstrating that the companion to NLTT~11748 is a cold ($T_{\\rm eff} \\lesssim 7400$\\,K), old (1.5--3\\,Gyr) C/O WD. We show in Section \\ref{sec:res} that eclipse modeling yields the first radius measurement of an extremely low-mass WD.  We close in Section \\ref{sec:conc} by discussing the value of additional measurements for constraining WD mass--radius relations.  \\begin{figure} \\centering \\epsscale{1.0} \\plotone{f1.eps} \\caption{Top panel: phased radial velocities from Keck/HIRES H$\\alpha$ measurements over two consecutive nights.  The black line is the ($e=0$) sinusoid with $K_1=271 \\, \\kms$ amplitude and the $v_r = 133 \\, \\kms$ systemic velocity offset.  Inset shows the phase combined H$\\alpha$ spectrum, in arbitrary flux units, with fitting function over plotted.  Bottom panel: phased photometric light curve for all three nights of data.  The error bar represents $4-5$\\,mmag uncertainty.} \\label{fig:rvlc} \\end{figure}   \\begin{deluxetable}{l c c c} \\tablecaption{Measured and Derived Quantities for NLTT~11748\\label{tbl:params}} \\tablehead{ \\colhead{Quantity} & \\mc{3}{c}{Value}\\\\ \\tableline \\mc{4}{c}{HIRES spectra} } \\startdata Rad.\\ vel.\\ amp., $K_1$ (km\\,s$^{-1}$)  & \\mc{3}{c}{271(3)} \\\\ Sys. radial velocity, $v_{r}$ (km\\,s$^{-1}$)  & \\mc{3}{c}{133(2)} \\\\ $\\chi^2$/dof         & \\mc{3}{c}{4.5/6} \\\\ \\cutinhead{FTN photometry} Time of prim.\\ ecl. (BJD TDB)  & \\mc{3}{c}{2,455,196.87828(7)} \\\\ Period (days)  & \\mc{3}{c}{0.2350606(11)} \\\\ Ephemeris $\\chi^2$/dof         & \\mc{3}{c}{3.5/3} \\\\ Primary eclipse depth, $d_1$  & \\mc{3}{c}{0.067(3)} \\\\ Secondary eclipse depth, $d_2$  & \\mc{3}{c}{0.034(2)} \\\\ Primary ecl.\\ duration, $\\tau_1$ (s) & \\mc{3}{c}{180(6)\\phz} \\\\ Secondary ecl.\\ duration, $\\tau_2$ (s) & \\mc{3}{c}{185(10)} \\\\ $F_2/F_1$ (in SDSS-$g^\\prime$)  & \\mc{3}{c}{0.035(3)} \\\\ Out of eclipse\\tablenotemark{a} $\\chi^2$/dof         & \\mc{3}{c}{404.8/372} \\\\ \\cutinhead{Combined data (assuming Mass of Primary $M_1=0.15M_{\\odot}$)} Limb darkening coefficient, $u_{\\rm LD}$  &  0.0  &  0.3  &  0.5  \\\\ Mass of secondary, $M_2$ ($M_{\\odot}$)   & 0.71(2)  &  0.71(2)  &  0.71(2)\\\\ Inclination (deg)  & 89.90(11) & 89.88(11) & 89.87(11)\\\\ Radius of primary, $R_1$ ($R_{\\odot}$)   & 0.0393(9) & 0.0406(9) & 0.0415(9)\\\\ $\\chi^2$/dof\\tablenotemark{b}         & 285.1/227 & 279.5/227 & 276.5/227\\\\ Distance\\tablenotemark{c} (pc)  & \\mc{3}{c}{150(32)} \\\\ Sys.\\ kin.\\ $(U,\\,V,\\,W)$\\tablenotemark{d} (km\\,s$^{-1}$)  & \\mc{3}{c}{(-151(9), -183(41), -34(5))} \\\\ \\enddata \\tablecomments{Quantities in parentheses are 1$\\sigma$ uncertainties on the last digit.} \\tablenotetext{a}{Excluding $\\pm0.005$ cycles around each eclipse.} \\tablenotetext{b}{Result of fitting phases $\\pm 0.1$ cycles around each eclipse.} \\tablenotetext{c}{Scaled from \\citet{kaw09} using our $R_1$.} \\tablenotetext{d}{Calculated from the proper motion \\citep{kaw09} corrected to the local standard of rest \\citep{hog05} and our updated distance.} \\end{deluxetable} ",
        "conclusions": "\\label{sec:conc}  Our discovery of the first eclipsing detached double WD binary has substantially constrained low-mass He WD models with radius measurements that agree with current models \\citep{ser02,pan07}. Additional measurements will reveal even more about the faint $\\approx 0.7 \\, M_{\\odot}$ C/O WD secondary. With our discovery, there are 11 confirmed double WD binaries that will merge within a Hubble time. From spectroscopic masses alone, six of these contain low-mass ($<0.25M_\\odot$) He core primaries with large ($>0.03 \\, R_\\odot$) radii indicative of a stable H burning shell \\citep{mul09,kil09,bad09,mar10,kul10}. Other than its inclination, NLTT~11748 is not extraordinary in any way relative to these, many of which offer larger solid angles for eclipse observations than NLTT~11748.  Our analysis of this system is far from complete, but our existing constraints, Table \\ref{tbl:params} and Figure \\ref{fig:rpvsmp}, highlight the value of further studies. High cadence photometry, resolving the ingress/egress of each eclipse, would yield the component radii as functions of inclination and orbital semimajor axis. If future spectroscopic observations can reveal the secondary's line features and measure its radial velocity, this would allow precise measurements of the masses of both components.  Detection of eclipses in different filters, both bluer and into the infrared, would additionally constrain the temperature of the secondary further as well as its age.  Model-independent mass and radius determinations for He and C/O WDs are rare.  For He WDs, there exist several pulsar binary systems for which Shapiro time delay measurements have yielded accurate mass determinations: \\object[PSR B1855+09]{PSR~B1855+09} \\citep{kas94}, \\object[PSR J0437-4715]{PSR~J0437$-$4715} \\citep{van01}, \\object[PSR J0751+1807]{PSR~J0751+1807} \\citep{bas06b}, \\object[PSR J1012+5307]{PSR~J1012+5307} \\citep{cal98,van96}, \\object[PSR J1911-5958A]{PSR~J1911$-$5958A} \\citep{bas06a}, and \\object[PSR J1909-3744]{PSR~J1909$-$3744} \\citep{jac03,jac05}. Of these, only four have radii measurements, PSR~J0437$-$4715, PSR~J0751+1807, PSR~J1911$-$5958A and PSR~J1909$-$3744.  However, all rely upon either distance measurements and evolutionary model-dependent luminosities or gravity measurements derived from atmospheric models.  Therefore, He WD evolutionary models \\citep{alt98,ser02,pan07} remain largely unconstrained. For C/O WDs, mass and radius measurements have been more frequent \\citep{sch96,pro98,pro02,cas09} and all rely on distance measurements via parallax or  cluster membership, surface brightness measurements and detailed knowledge of each spectrum.  To reach the precisions required to constrain theoretical models, these measurements must rely upon atmospheric model spectra.  Only recently \\citet{par10} have derived precise measurements of mass and radius in a model-independent way in the eclipsing binary system \\object[NN Ser]{NN~Ser}.  NLTT~11748 offers an additional C/O WD system for model-independent mass and radius measurements as well as the first such system for He WDs.  Gravitational wave emission will bring this system into contact in $13.8-6.3$\\,Gyr (for $M_1=0.1-0.2 \\, M_\\odot$).  With a mass ratio of $0.15/0.71=0.2$, stable mass transfer at contact is likely to occur and create an AM~CVn binary \\citep{nel01,mar04}.  The known orbital period, inclination, and distance ($d\\approx150$\\,pc) yield the gravitational wave strain at Earth, $h=(3.6-7.9)\\times 10^{-23}$ \\citep{tim06,roe07} at a frequency $\\nu\\approx 10^{-4}$\\,Hz for $M_1=0.1-0.2 \\, M_\\odot$. Though not `louder' than the verification sources tabulated by \\citet{nel09}, NLTT~11748's accurate ephemeris allows for a coherent folding of the {\\it LISA} data over the mission duration."
    },
    "1005/1005.2286_arXiv.txt": {
        "abstract": "We derive a generic weak lensing equation and apply it for the study of images produced by tidal charged brane black holes. We discuss the similarities and point out the differences with respect to the Schwarzschild black hole weak lensing, to both first and second order accuracy, when either the mass or the tidal charge dominates. In the case of mass dominated weak lensing, we analyze the position of the images, the magnification factors and the flux ratio, as compared to the Schwarzschild lensing. The most striking modification appears in the flux ratio. When the tidal charge represents the dominating lensing effect, the number and orientation of the images with respect to the optical axis resembles the lensing properties of a Schwarzschild geometry, where the sign associated with the mass is opposite to that for the tidal charge. Finally it is found that the ratio of the brightness of images as a function of image separation in the case of tidal charged black holes obeys a power-law relation significantly different from that for Schwarzschild black holes. This might provide a means for determining the underlying spacetime structure. ",
        "introduction": "Gravitational lensing has become a useful tool in measuring certain properties of gravitational fields ever since the beginnings of general relativity. While the initial observations of gravitational lensing were used to verify the theoretical predictions of general relativity, it has now been employed to study the large scale structure of the Universe, to determine behaviour of compact stellar objects and to search for dark matter candidates. In what follows we propose the idea that gravitational lensing might also be used to determine which among various gravitational theories is correct. It is already well known that the predictions of bending angles computed from Newtonian gravity compared to those obtained from Einstein's gravity differ by a factor of two. Therefore using gravitational lensing a means for exploring the differences between competing gravitational theories may well provide a technique that can be employed to determine the dimensionality of spacetime, or the coupling of matter and fields to the gravitational field or even to distinguishing among different formulations of gravity theories.  In general the path taken by photons in a gravitational field provide a number of different effects that can be measured using current state-of-the-art telescopes. The first and most obvious one is the production of multiple images and the relative separation of those images. Secondly gravitational lensing can produce a change in brightness of the images depending on how much bending a group of initially parallel rays undergoes. Finally if the source or lensing object has a time dependency, the changes in arrival times of light signals can provide a very accurate measurement of some spacetime properties.  In this paper, we discuss the formation of images (i.e. their location and brightness) for some black holes that are predicted to exist in 5-dimensional brane-world theories. By computing the bending angles and image brightness changes that can occur due to the passage of photons past the objects described by black hole solutions of the theory, we should be able to determine enough of the properties of the lensing object to distinguish a general relativistic black hole from that predicted by an alternative theory.  Brane-world models have standard-model matter confined to a 3+1 dimensional hypersurface, and gravity acting in a higher-dimensional non-compact space-time. Such models have attracted much attention in recent years, both as candidate theories meant to solve the hierarchy problem and predicting modified cosmological evolution \\cite{BDEL}. Higher codimension branes were considered in connection with conical singularities \\cite{BGNS}, \\cite% {deRham}. For a codimension-one brane-world, matter is generated on the brane by the junction conditions though the brane representing a discontinuity in the extrinsic curvature \\cite{MaartensLivRev}- \\cite% {VarBraneTensionPRD}. Although the early expectation to replace dark energy was not met, codimension-one brane-worlds still can produce alternative explanations for dark matter \\cite{MakHarko}-\\cite{Pal}. Therefore the study of localized matter configurations, in particular black hole solutions admitted in brane-world theories became important.  Analytic black hole solutions include six-dimensional locally Schwarzschild solutions \\cite{Kaloper}; static five-dimensional black holes localized on the brane, with the horizon decaying in the extra dimension and generated by energy condition violating shells \\cite{KantiTamvakis}, \\cite% {KantiOlasagastiTamvakis}; and black holes with a radiating component in the extra dimension \\cite{EmparanFabbriKaloper}. Numerically small brane black holes (compared to the five-dimensional curvature) were shown to exist as five-dimensional Schwarzschild solutions \\cite{KudohTanakaNakamura} in asymptotically five-dimensional Anti de Sitter space-time.  The four-dimensional Schwarzschild metric can be also extended into the fifth dimension as a black string \\cite{ChRH}, which due to the Gregory-Laflamme instability \\cite{GL} can pinch off, leading to a black cigar metric \\cite{Gregory} (although under very mild assumptions, classical event horizons will not pinch off \\cite{HorowitzMaeda}). Gravity wave perturbations of such a black-string brane-world were computed in Ref. \\cite% {SCMlet}.  The perturbative analysis of the gravitational field of a spherically symmetric source in the weak field limit in the original Randall-Sundrum setup (Schwarzschild black hole on a brane embedded in Anti de Sitter five-dimensional space-time) has shown corrections to the Schwarzschild potential scaling as $r^{-3}$ \\cite{RS}, \\cite{GarrigaTanaka} and \\cite% {Giddings}. However if the Schwarzschild black hole is embedded in another higher-dimensional space-time, this scaling would not apply \\cite% {BlackString}. Both weak \\cite{KarSinha}, \\cite{MM} and strong \\cite{Whisker} gravitational lensing of various brane black holes were discussed, the topic being reviewed in Ref. \\cite{MajumdarMukherjee}.  The effective Einstein equation on the codimension-one brane admits a spherically symmetric vacuum solution characterized by two parameters: the mass $m$ and a tidal charge $q$, the latter arising from the Weyl curvature of the 5-dimensional space-time in which the brane is embedded \\cite{tidalRN}% :% \\begin{eqnarray} ds^{2} &=&-f\\left( r\\right) dt^{2}+f^{-1}\\left( r\\right) dr^{2}+r^{2}\\left( d\\theta ^{2}+\\sin ^{2}\\theta d\\varphi ^{2}\\right) ~,  \\nonumber \\\\ \\ f &=&1-\\frac{2m}{r}+\\frac{q}{r^{2}}~.  \\label{tidal} \\end{eqnarray}% For $q\\leq m^{2}$ this represents a black hole with horizons given by $% r_{\\pm }=m\\pm \\left( m^{2}-q\\right) ^{1/2}$. For $q<0$ only $r_{+}\\,$\\ is positive, therefore there is only one horizon. For $q=0$ the line element describes the Schwarzschild metric and for $0<q\\leq m^{2}$ it is formally identical to the general relativistic Reissner-Nordstr\\\"{o}m electro-vacuum solution with electric charge $Q=q^{1/2}$. In the limit $q=m^{2}$ the metric becomes extremal, such that the two horizons coincide. For $q>m^{2}$ there is no horizon at all. The metric is singular at $r=0$, thus it describes a naked singularity. It should be noted that a negative tidal charge strengthens gravity (the horizon is outside the Schwarzschild radius $2m$), such that $q<0$ contributes to the localization of gravity on the brane. A positive tidal charge weakens gravity, both horizons lying below the Schwarzschild radius for $0<q\\leq m^{2}$, and obstructing the apparition of a horizon at all for $q>m^{2}$.  Observations on light deflection could in principle constrain both the lensing black hole parameters and the underlying gravitational theory. The lensing properties of a Schwarzschild geometry were thoroughly investigated in Refs. \\cite{VE}, \\cite{Virbhadra}. Recently in Ref. \\cite{tidalDeflection} the deflection angle of light rays passing near the tidal-charged brane black hole was computed up to the second order in the perturbation theory. The Hamiltonian method gave identical results to the previously employed Lagrangian method \\cite{PADEU}. The light deflection was derived in terms of the small parameters $\\varepsilon :=m/b$ and $\\eta :=q/b^{2}$. Here $b$\\ is the impact parameter, defined as the distance of the lensing object to the straight line trajectory, which would occur in the absence of the lensing object. The deflection angle to second order accuracy is \\begin{equation} \\delta =4\\varepsilon -\\frac{3}{4}\\pi \\eta +\\frac{15}{4}\\pi \\varepsilon ^{2}-16\\varepsilon \\eta +\\frac{105}{64}\\pi \\eta ^{2}~.  \\label{delta} \\end{equation}% A first confrontation with Solar System measurements in Refs. \\cite% {tidalDeflection}, \\cite{BoehmerHarkoLobo} led to constraints on $q$ and on the brane tension $\\lambda $.  Further exploring the consequences of the result (\\ref{delta}), in the present paper we study the formation of images and their magnification factors in the tidal charged black hole geometry, focusing on the similarities and differences with the purely general relativistic case.  In the case of weak lensing where the latter black hole solutions differ from standard general relativistic black holes, we are able to develop a perturbative computation of the bending angles and image brightnesses that can be expected when the black hole masses produce the dominant effect. However the restriction to objects with weak tidal charge (compared to mass) is not necessary. As long as the impact parameter or distance of closest approach for the light ray remains large compared to a natural length scale associated with the tidal charge, the weak field limit is maintained and tidally charged black hole light bending effects can also be computed.  In Section 2 we study the weak lensing under quite generic circumstances. We derive a lens equation for weak lensing, which will allow for a study of the lensing and image formation to second order accuracy in the chosen small parameters. This lens equation is more generic than, and reduces to, the Virbhadra-Ellis equation \\cite{Virbhadra}-\\cite{Bozza} in a properly defined approximation. We give the explicit form of both lens equations applied to the tidal charged black hole in the Appendix, in order to see at which order the differences arise.  In Section 3 we employ the deflection angle (\\ref{delta}) in our lens equation and obtain an approximate equation for weak lensing of the tidal charged black hole in a form of a cubic polynomial. For later comparison we also review here the lensing by Schwarzschild black holes, which arises in our formalism to first order when the mass parameter dominates over the tidal charge parameter.  In Section 4 we discuss the corrections in the location of the images and magnification factors, by including the contributions quadratic in the mass parameter and linear in the tidal charge parameter. This analysis leads to results that are similar to the general relativistic Reissner-Nordstr\\\"{o}m black hole lensing, some of them discussed in Ref. \\cite{Sereno}. Despite $q$ being similar to the square of the electric charge of the Reissner-Nordstr% \\\"{o}m black hole, the case of a negative tidal charge is without counterpart in general relativity.  Section 5 contains the analysis of the case, where the tidal charge dominates over the mass, by keeping only the first order tidal charge contribution. We determine the image locations and the magnification factors. In Section 6 we discuss the second order corrections to the above case, by including contributions linear in the mass and quadratic in the tidal charge.  We summarize our findings in the Concluding Remarks. Our choice of units is given by $G=1=c$. ",
        "conclusions": "Using simple geometric relations we derived a generic lens equation, Eq. (% \\ref{LE}) for weak lensing. This formula is more accurate than the Virbhadra-Ellis lens equation (as it contains no approximations of trigonometric expressions), but reduces to it in a proper limit (differences are to be expected in asymmetric source and observer distances with respect to the lens).  We have applied our lens equation in the discussion of the weak lensing by tidal charged black holes, to both first and second order accuracy in the black hole parameters, when either the mass or the tidal charge dominates. We have carried on expansions in the small mass and tidal charge parameters, as well as in the angles spanned by the real and apparent positions of the sources with the optical axis. In the Appendix we have investigated the differences between the predictions of the two lens equations. The predictions of our lens equation (\\ref{LE}) and of the Virbhadra-Ellis lens equation (\\ref{lens_Ellis}) coincide in most of the cases we consider, with the notable exception of the tidal charge dominated black hole lensing discussed to a higher accuracy in Section \\ref{2o_q}, where the Virbhadra-Ellis lens equation would not predict (or would predict with different coefficients) some of the higher order terms.  Although Solar System tests lead to the expected result, that light deflection by the Sun is due to its mass, and leaves room but for a minuscule tidal charge \\cite{tidalDeflection}, \\cite{BoehmerHarkoLobo}, it cannot be excluded that tidal charge dominated black holes could exist on the brane, as the tidal charge is an imprint of the Weyl curvature of the higher-dimensional space-time, which remains unspecified for the tidal charged black hole.  In the case of mass dominated weak lensing, we found that the position of the images is similar to the Reissner-Nordstr\\\"{o}m black hole lensing, discussed in Ref. \\cite{Sereno}. We have analyzed, how the image separations, the magnification factors and the flux ratios are modified as compared to the Schwarzschild lensing by the perturbations arising from second order mass and linear tidal charge contributions ($\\bar{\\varepsilon}% ^{2}$ and $\\bar{\\eta}$, respectively). The most striking modification appears in the ratio of the magnification factors (the flux ratio), shown on Fig \\ref{mag1}, which can be either increased or decreased, depending on the sign of $\\bar{\\eta}-5\\bar{\\varepsilon}^{2}$.  When the tidal charge dominates in the lensing behaviour, the situation is different. The case of positive tidal charge resembles the lensing properties of a negative mass Schwarzschild spacetime \\cite{Cramer}.  Black holes with negative tidal charge are however favoured by strengthening and confining gravity to the brane and also by thermodynamic considerations \\cite{tidalThermo}. In the case of a dominant negative tidal charge the lensing properties are similar to those of a positive mass Schwarzschild black hole, where the similarity is only in the number of images lying above or below the optical axis. The actual location of the images is different and this fact is summarized in Fig \\ref{fig3}, which is one of the main results of this paper.  Finally the power law dependence of the ratio of the magnification factors on the separation of the images provides an means for observing the differences between the Schwarzschild and tidal charged black holes. Given that the next generation of radio telescopes will easily be able to resolve images to less than milli-arcsecond accuracy, the different rates at which the ratio of the brightness changes should be able to provide a significant observational signature to constrain the Weyl curvature as a substitute for dark matter."
    },
    "1005/1005.2415_arXiv.txt": {
        "abstract": "s{We use a variant of principal component analysis to investigate the possible temporal evolution of the dark energy equation of state, $w(z)$. We constrain $w(z)$ in multiple redshift bins, utilizing the most recent data from Type Ia supernovae, the cosmic microwave background, baryon acoustic oscillations, the integrated Sachs-Wolfe effect, galaxy clustering, and weak lensing data. Unlike other recent analyses, we find no significant evidence for evolving dark energy; the data remains completely consistent with a cosmological constant. We also study the extent to which the time-evolution of the equation of state would be constrained by a combination of current and future-generation surveys, such as Planck and the Joint Dark Energy Mission.} ",
        "introduction": "One of the defining challenges for modern cosmology is understanding the physical mechanism responsible for the accelerating expansion of the Universe~\\cite{Riess:1998cb,Perlmutter:1998np}. The origin of the cosmic acceleration can be due to a new source of stress-energy, called ``dark energy'', a modified theory of gravity, or some mixture of both \\cite{Uzan:2006mf}.\\\\ In the absence of a well-defined and theoretically motivated model for dark energy, it is generally assumed that the dark energy equation of state (the ratio of pressure to energy density) evolves with redshift with an arbitrary functional form. Common parameterizations include a linear variation, $w(z)=w_0+w_zz$~\\cite{Cooray1999}, or an evolution that asymptotes to a constant $w$ at high redshift, $w(z)=w_0+w_az/(1+z)$~\\cite{Chevallier2001,Linder2003}. However, given our complete ignorance of the underlying physical processes, it is advisable to approach our analysis of dark energy with a minimum of assumptions. Fixing an ad hoc two parameter form could lead to bias in our inference of the dark energy properties.\\\\ In this paper we measure the evolution history of the dark energy using a flexible and almost completely model independent approach, based on a variant of the principal component analysis (PCA) introduced in Huterer (2003)~\\cite{Huterer2003}; in order to be conservative, we begin by using data we determine the equation only from geometric probes of dark energy, namely the cosmic microwave background radiation (CMB), Type Ia supernovae (SNe) and baryon acoustic oscillation data (BAO). We perform a full likelihood analysis using the Markov Chain Monte Carlo approach~\\cite{Lewis2002}. We then consider constraints on $w(z)$ from a larger combination of datasets, including probes of the growth of cosmological perturbations, such as large scale structure (LSS) data. An important consideration for such an analysis is to properly take into account dark energy perturbations, and we make use of the prescription introduced in~\\cite{Feng2008}. This method implements a Parameterized Post-Friedmann (PPF) prescription for the dark energy perturbations following~\\cite{Hu2007}. ",
        "conclusions": "One of the main tasks for present and future dark energy surveys is to determine whether or not the dark energy density is evolving with time.\\\\ We have performed a global analysis of the latest cosmological datasets and have constrained the dark energy equation of state using a very flexible and almost model independent parameterization. We determine the equation of state $w(z)$ in five independent redshift bins, incorporating the effects of dark energy perturbations. We find no evidence for a temporal evolution of dark energy---{\\em the data is completely consistent with a cosmological constant}. This agrees with most previous results, but significantly improves the overall constraints. We show that future experiments, such as Planck or JDEM, will be able to reduce the uncertainty on $w(z)$ to less than $10\\%$ in multiple redshift bins, thereby mapping any temporal evolution of dark energy with high precision. With this data it will be possible to measure the temporal derivative of the equation of state parameters, ${dw}/{dz}$, useful in discriminating between two broad classes of ``thawing'' and ``freezing'' models~\\cite{Caldwell2009}."
    },
    "1005/1005.0626_arXiv.txt": {
        "abstract": "Interacting binaries typically have separations in the milli-arcsecond regime and hence it has been challenging to resolve them at any wavelength. However, recent advances in optical interferometry have improved our ability to discern the components in these systems and have now enabled the direct determination of physical parameters. We used the Navy Prototype Optical Interferometer to produce for the first time images resolving all three components in the well-known Algol triple system. Specifically, we have separated the tertiary component from the binary and simultaneously resolved the eclipsing binary pair, which represents the nearest and brightest eclipsing binary in the sky. We present revised orbital elements for the triple system, and we have rectified the 180-degree ambiguity in the position angle of Algol C. Our directly determined magnitude differences and masses for this triple star system are consistent with earlier light curve modeling results. ",
        "introduction": "Algol ($\\beta$ Per, HR 936), the prototype for a well known class of eclipsing binaries, has well over 200 years of published observations available for study. The eclipsing nature of the system was first suggested by \\citet{good} when he stated the light variations could result from ``... the interposition of a large body revolving around Algol ...''. A long suspected third component in Algol was spectroscopically confirmed by \\citet{ss57} and \\citet{ebb}. \\citet{soder} provides a review of Algol which is a useful starting point for summarizing this triple system. We will refer to the B-type primary star in Algol as Algol A, the K type secondary star as Algol B, and the more distant Am companion as Algol C \\citep[see][for the exact spectral types]{mtr93}. It is thus a hierarchical triple system as defined by \\citet{evans}. \\citet{wh72} made the first radio detection of Algol and \\citet{les93} determined that the radio emission in Algol comes from Algol B using multi-epoch Very Long Baseline Interferometry observations. Algol was one of 12 radio stars used to link the HIPPARCOS optical reference frame to the International Celestial Reference System (ICRS) \\citep{kov}. The 1.86 year orbit of Algol C was resolved by speckle interferometry \\citep{lab74} and optical interferometry \\citep{pan} but these investigations could not resolve the close binary. Also, these results suffer from a 180\\arcdeg\\ ambiguity in the absolute position angle (PA) due to a lack of phase information \\citep{lab70}. The Fourth Catalog of Interferometric Measurements of Binary Stars has a more complete listing of these observations\\footnote{http://www.usno.navy.mil/USNO/astrometry/ optical-IR-prod/wds/int4 \\citep{int}}.  This considerable body of knowledge still leaves room for additional exploration, and permitted inconsistent descriptions of the orbital elements of the triple system. It would appear that the orbital elements listed in \\citet{soder} and \\citet{pan} would adequately describe the Algol triple system. But these elements include a 180\\arcdeg\\ difference in the position angle of the ascending node when compared with the orbital solution in the Hipparcos catalog \\citep{linde}. Algol's role as one of a small number of radio stars used to link the HIPPARCOS optical reference frame to the ICRF requires resolution of this inconsistency. Algol can also serve as a position angle calibrator for optical and near-IR interferometers and this strengthens the case for a resolution of this disagreement.  Recently, Algol was observed in the near IR (K$_{s}$ = 2.133 $\\mu$m) with the CHARA array and at 5 GHz with the European VLBI Network during December 2006 \\citep{csiz}. The CHARA array with approximately 200 m baselines has a similar spatial resolution to the NPOI at optical wavelengths with 64 m baselines. The CHARA data resolve the close binary but \\citet{csiz} make no mention of detecting Algol C. Csizmadia et al. use their VLBI and CHARA array observations to produce an orbit of Algol A$-$B with an opposite sense of rotation from that determined by \\citet{les93}. We are then presented with another inconsistency in published results for the Algol system.  During 2006 October and November we collected observations of Algol using the Navy Prototype Optical Interferometer (NPOI), the Very Long Baseline Array (VLBA)\\footnote {The National Radio Astronomy Observatory is operated by Associated Universities, Inc., under cooperative agreement with the National Science Foundation.}, and the Lowell Observatory 42'' Hall telescope equipped with the solar$-$stellar spectrograph. We defer discussion of the radio and spectroscopic results for a future paper. The primary aim of this project was to resolve the close pair in Algol with the NPOI, and perform an absolute astrometric registration of the optical NPOI images to the ICRF phase-referenced VLBA images. As the NPOI records visibility squared and closure phase data we can determine the position angle calibration without a 180\\arcdeg\\ ambiguity and resolve the inconsistency between the orbital elements in \\citet{soder} and \\citet{pan} and the orbital elements in the HIPPARCOS catalog. We can also address the inconsistency in the direction of the close binary orbit between \\citet{les93} and \\citet{csiz}.     In this paper, we report on our NPOI observations which extend the direct knowledge of the Algol triple system to optical wavelengths. In \\S2 we discuss our observations, with particular attention to the calibration of the NPOI absolute position angles. In \\S3 we present the astrometric orbits of the Algol A$-$B and AB$-$C systems. We conclude with \\S4 and a comparison with light curve solutions and a discussion of the astrometric results. ",
        "conclusions": "Our discussion of the analysis of the interferometric observations will be restricted to the new insights we have gained on this well studied system. Our NPOI observations mark the first resolution of the Algol system into three components. \\citet{csiz} reported the resolution of the close binary with the CHARA array at near infrared wavelengths but without absolute phase calibration. Our orbital solution is fully consistent with the pioneering radio interferometric observations of \\citet{les93}. We have unambiguously determined that the close pair orbit is retrograde, and nearly orthogonal to the plane of the wide orbit. The relative angle $\\phi$ is $96^{\\circ} \\pm 5^{\\circ}$ according to  \\[ \\cos \\phi = \\cos i_1\\cos i_2+\\sin i_1\\sin i_2\\cos(\\Delta \\Omega), \\]  \\noindent where  $i_1$ and $i_2$ are the inclinations of the two orbits, and $\\Delta\\Omega$ is the difference between the two ascending node angles. This improvement to the orbital plane orientation is relevant to dynamical studies of hierarchical triples \\citep[e.g.][]{kem98}. This orientation of the wide orbit also removes the discrepancy expected between the photocenter motion one would compute using the AB$-$C elements of \\citet{pan} and that found from HIPPARCOS or the orbit of \\citet{heintz94}. Our retrograde orientation of the close binary orbit contrasts with the prograde orbit of \\citet{csiz}. The position angle calibration of \\citet{csiz} depended on VLBI observations made during a radio flare of Algol which may have complicated their analysis due to the changing radio morphology of Algol \\citep[\\S 4.1 of ][]{csiz}. As the NPOI observations are calibrated to produce an absolute position angle as shown in \\S 2 (with $\\gamma$ Per as the position angle calibrator) we are confident that the retrograde orbit of the close binary pair is correct. The 15 GHz VLBI observations of \\citet{peter} also agree with the previously determined retrograde orbit.  The absolute position angle calibration of the NPOI enables a revision of the orientation of the Algol C orbit of \\citet{pan}. Our results place the maximum AB$-$C separation in the same quadrant as found by Heintz \\citep{heintz94,gjh}. The Hipparcos orbital solution \\citep[Double and Multiple Systems Annex]{esa97} references \\citet{gjh} for the quadrants of the longitude of periastron $\\omega$ and the position angle of the ascending node $\\Omega$, and our observations confirm the accuracy of the Hipparcos orbital elements. This confirmation is important because it links results from both the optical and radio reference frames. A 180$^{\\circ}$ reversal of the position angle of the ascending node of the AB$-$C orbit presented here would create a time variable systematic offset of the photocenter \\citep{gjh} that could not be reconciled with the HIPPARCOS observations.  Our interferometric observations have resulted in the first resolved images of the triple system and the first directly measured magnitude differences for the three stars in Algol. Previous estimates of the magnitude difference for the close binary were made by modeling photometric and spectroscopic data \\citep[e.g.][ and references therein]{kim,mtr88}. The V band magnitude differences predicted from these models span slightly more than one magnitude: 3.72 $\\pm$ 0.10 \\citep{w72}, 2.97 $\\pm$ 0.31 \\citep{soder}, 3.92 $\\pm$ 0.88 \\citep{mtr88}, and 2.71 $\\pm$ 0.15 \\citep{kim}. Our directly measured V band magnitude differences (Table~\\ref{tbl-3}) favor magnitude differences of less than 3. Our AB$-$C magnitude difference is in excellent agreement with the previously determined value using the Mark III \\citep{pan} and it is also consistent with an early speckle interferometry result \\citep{lab74}. We extrapolated the magnitude difference to the center wavelength of the Hipparcos Hp filter using the stellar atmosphere parameters from Table~\\ref{tbl-2}, and used the masses given in that table to determine a $18.4$ mas amplitude for the motion of the photo center. This result is consistent with the Hipparcos orbital solution of $19.0\\pm0.6$ mas (ESA 1997). The dynamical parallax determined from our full fit to the astrometric data and published radial velocities is $34.7\\pm0.6$, consistent with the Hipparcos value of $35.1\\pm0.9$ mas.  Our magnitude differences add a directly measured constraint to the results obtained from modeling the photometric light curve and spectroscopic data. It may be useful to re-examine the modeling of the close binary using magnitude differences derived directly from the interferometer measurements. Other bright double and multiple stellar systems will yield similar constraints for use with spectroscopic and photometric data."
    },
    "1005/1005.3139_arXiv.txt": {
        "abstract": "We present the first detailed structure formation and radiative transfer simulations of the reionization history of our cosmic neighbourhood. To this end, we follow the formation of the Local Group of galaxies and nearby clusters by means of constrained simulations, which use the available observational constraints to construct a representation of those structures which reproduces their actual positions and properties at the present time. We find that the reionization history of the Local Group is strongly dependent on the assumed photon production efficiencies of the ionizing sources, which are still poorly constrained. If sources are relatively efficient, i.e. the process is 'photon-rich', the Local Group is primarily ionized externally by the nearby clusters. Alternatively, if the sources are inefficient, i.e. reionization is 'photon-poor' the Local Group evolves largely isolated and reionizes itself. The mode of reionization, external vs. internal, has important implications for the evolution of our neighbourhood, in terms of e.g. its satellite galaxy populations and primordial stellar populations. This therefore provides an important avenue for understanding the young universe by detailed studies of our nearby structures. ",
        "introduction": "Approximately thirteen billion years ago the cosmic neighbourhood destined eventually to become our Local Group of galaxies underwent a dramatic transition: a giant ionization front swept through, engulfing it in a sea of ionizing radiation. This occurred as a local manifestation of a global transition of the intergalactic medium in the whole universe referred to as Cosmic Reionization, caused by the radiation from the first galaxies. This process converted the intergalactic medium from neutral and cold gas during the Cosmic Dark Ages before any galaxies existed, into a hot, ionized plasma.  The absorption spectra of QSOs from redshift 0 to about 6 show that the intergalactic medium has been almost fully ionized for most of the lifetime of the Universe. On the other hand, the recent data from the Wilkinson Microwave Anisotropy Probe (WMAP) satellite yielded a rather large optical depth for scattering the Cosmic Microwave Background photons on free electrons. This strongly suggests that the reionization epoch started well before redshift 10 and therefore was fairly extended in time. It also confirmed independently the existence of a reionization epoch, required by this additional optical depth. The process of reionization had far-reaching consequences for subsequent galaxy formation. The photoionization heating which accompanies reionization increased the gas temperature from the very low one, of order a few K or less, during the Cosmic Dark Ages before the first stars formed, to $\\sim10^4$~K or more. This in turn increased the corresponding Jeans mass, the mass above which gas pressure cannot sucessfully counteract gravity, by about 5 orders of magnitude. This strongly suppressed the formation of low-mass galaxies and cut off the star formation in previously-formed ones, and thereby should have significantly influenced the early population of dwarf satellite galaxies. For this reason, reionization is often invoked as a plausible explanation for the observed lack of galaxy satellites compared to the numbers predicted by pure dark matter simulations \\citep{2000ApJ...539..517B, 2009MNRAS.400.1593M, 2010MNRAS.402.1995M,2010ApJ...710..408B}. The same ionization and heating process also changed significantly the character of star formation, as stars whose formation starts out from hot, ionized gas, even those with a primordial element abundance, are different from the ones forming out of initially cold gas \\citep[e.g.][]{2003ApJ...586....1M}.  In recent years a variety of theoretical and numerical modelling work has shown that reionization was highly patchy in nature, with a very large variations in the times at which different regions become reionized \\citep[e.g.][]{2006MNRAS.369.1625I, 2007MNRAS.376..534I,2007ApJ...654...12Z}. Many of the effects of reionization on the later structure formation are highly dependent on the stage at which the forming structures are at the time when they become reionized. Therefore, it is important to follow the reionization history of different patches in detail and to compare the timing of reionization with the stage of formation of the local structures.  \\citet{2007MNRAS.381..367W} used large-scale reionization simulations to study the reionization history of galaxies of a variety of present-day types: field galaxies, cD central cluster galaxies, $L_*$ galaxies and Local Group-like systems. In particular, they focused on the question which galaxies are reionized internally (i.e. by their own progenitors) vs. externally (i.e. before sigificant fraction of their mass is in collapsed objects). For most galaxy types the answer is statistical, with certain probabilities for each outcome. They found that there is a halo mass scale, of order $10^{12}M\\odot$, which divides the massive galaxies which are predominantly internally reionized and lower-mass ones for which the opposite is true. More recently, \\citet{2009ApJ...703L.167A} studied the same problem using larger-scale, coarsely-resolved semi-analytical reionization calculations. They achieved better statistics for the larger halos, from Milky Way-sized up to galaxy clusters, as a consequence of the larger volume their calculation followed and reached very similar conclusions to \\citet{2007MNRAS.381..367W}.  However, these previous studies are statistical and therefore can only yield a certain probability for a given type of system to be externally or internally reionized. In contrast, the answer for a specific system, like our own Local Group of galaxies, consisting of the Milky Way, Andromeda, M33 and their satellite galaxies, will depend on the details of its and its neighbouring systems' formation, the timing of that formation, and the relative positions in space of their projenitors at the relevant epochs. This kind of detailed information can only be obtained through numerical simulations with constrained initial conditions.  Constrained simulations aim to reproduce the spatial and velocity structure of our Local Group and its neighbourhood at present (redshift $z=0$). In this work we combine the most advanced constrained realizations of the Local Group available at present with an accurate treatment of the radiative transfer during cosmic reionization. This allowed us for the first time to calculate the specific reionization history of all Local Group progenitors, as well as those of the nearby structures like the Virgo and Fornax clusters of galaxies. ",
        "conclusions": "We have performed the first simulations of the reionization history of our local neighbourhood of the universe based on constrained N-body simulations of the formation of local structures - our Local Group of galaxies and the nearby galaxy clusters. The reionization history of the Local Group cannot yet be predicted uniquely, primarily due to the still poor observational constraints on the properties of the reionization sources. Therefore, we studied two models constructed so as to roughly bracket the range of expected outcomes. While both models satisfy the available global observational constraints, they differ significantly in their underlying assumptions. Our first model assumes relatively high ionizing photon production efficiencies ('photon-rich reionization'), while our second model studies the other extreme, where the sources produce barely enough photons to complete the reionization process in time ('photon-poor reionization'). The second model therefore results in more extended reionization history and delayed final overlap. The two models have the same observational constraints imposed on their initial matter distribution, which guarantees that the  large-scale structures closely resemble our local neighbourhood, but they differ in their random component, yielding different constrained realizations. This allows us to check the robustness of our results with respect to the specific realization. Finally, the two models also have different underlying cosmological models (WMAP 3-year vs. WMAP 5-year best fit). However, the effect of the background cosmology is well understood and results in an overall shift of the reionization earlier or later, with no significant effects on our results, for which only the relative timing of structure formation vs. reionization history is of importance.  Our results show that the assumed efficiency of the ionizing sources has the most important influence on the nature of the reionization history of our Local Group of galaxies. Efficient photon production ensures that the nearby clusters emit more than sufficient number to ionize both themselves and their surroundings, including the Local Group. The fact that those galaxy clusters (Virgo and Fornax) coincide with high, rare peaks of the density field means that they form their progenitor halos earlier than the LG, which is in a more average region of the universe. As a result, the large-scale ionization fronts which propagated outward from the proto-clusters overrun the Local Group before it managed to form enough sources to ionize itself, resulting in its reionization being mostly externally-driven.  Several points are worth noting here. Although generally the radiative transfer is a highly non-local phenomenon, which feature complicates its numerical treatment and the code parallelization, during most of the EoR the situation is somewhat more complicated. The neutral patches have enormous optical depth to soft ionizing radiation (the only type of radiation we consider here). Even the already-ionized patches still have considerable continuum optical depth over cosmological (multiple Mpc) distances due to the small residual neutral fraction still remaining in such regions\\footnote{E.g. At redshift $z=8$, mean density and residual neutral fraction of $10^{-3}$ the mean free path is $\\sim3$~Mpc. Density fluctuations further diminish this value.}. This residual neutral fraction diminishes over time, but does so only gradually, as more and more sources appear and the mean flux thereby increases. As a consequence of all this, reionization starts out as a fairly local process where only the relatively nearby, directly visible ionizing sources within the same ionized bubble contribute to the flux at a given point. This property allows us to focus our analysis on the important local sources and ignore the far-away ones for our current purposes (they are of course all included in the radiative transfer simulation). In our $\\sim100$~Mpc box there are multiple proto-clusters which collapse nonlinearly by the present, but of those only Virgo and Fornax are sufficiently close to potentially contribute to the reionization of our Local Group.  Furthermore, the ionization fronts propagate through underdense regions (voids) much faster than through overdense ones (filaments, knots). Therefore, the relative positioning of the structures of interest and the density fluctuations in their immediate neighbourhood are important. Once the available observational constraints are imposed in order to reproduce the local structures, we find that the Local Group is separated from Virgo and Fornax by voids in either realization (see Figure~\\ref{dens_slice_fig}). In contrast, the previous studies of this problem which did not use constrained realizations \\citep{2007MNRAS.381..367W,2009ApJ...703L.167A} sampled a wide range of environments and relative positions of nearby clusters. Such, purely statistical approach yields valuable insights on the range of reionization histories that could be expected for a certain type object (e.g. LG-like objects). However, by its nature such approach necessarily includes many objects which, although they share certain basic features, locally do not reproduce the specific large-scale structures around us. Therefore, the constrained realizations are indispensible if we want to make realiable predictions for the effects of reionization on our neighbourhood.  Why is the mode of reionization, external vs. internal, of our Local Group an important issue? This has a number of important implications for the formation of structures. Reionization dramatically rises the Jeans mass, thus impeding the formation and growth of small galaxies. In terms of this effect, the galactic haloes fall into three categories. The gas in the smallest halos (minihalos), whose virial temperatures are below the limit ($\\sim10^4$~K) for efficient radiative cooling through atomic line radiation. The ionization of the gas brings its temperature to $\\sim10^4$~K and it boils out, resulting in their complete evaporation \\citep{2004MNRAS.348..753S,2005MNRAS...361..405I}, which leaves behind dark halos. In the other limit, the galaxies above certain mass ($M\\gtrsim10^{10}M_\\odot$) have sufficiently deep gravitational potential wells to sucessfully withstand the effects of ionizing radiation and are thus not significantly affected by the reionization process. The effects of radiative feedback on dwarf galaxies of intermediate mass, roughly between $10^8M_\\odot$ and $10^{10}M_\\odot$ is more complex and still very much a subject of active investigation. The gas in such already-formed systems cannot be photoevaporated, as it can cool back down to $\\sim10^4$~K very efficiently. However, photoionization heating rises the intergalactic gas temperature and pressure, which rises the Jeans mass and thereby suppresses the future formation of very low-mass galaxies, as well as curtails the fresh gas infall onto such halos. Larger galaxies are less affected directly, but could do so indirectly, through their smaller progenitors, which could be expected e.g. to result in smoother gas sub-structure and modified stellar populations. Where the boundary between efficient and inefficient feedback from reionization lies is still unclear and very much subject of active research. Full investigation of the effects of reionization on galaxy formation and satellite galaxy populations goes well beyond the scope of the current work. However, our present results indicate that the photon production efficiencies of the first galaxies are the main factor determining the type of reionization history which our Local Group underwent. Therefore, this process should have left useful fossil records in the properties of our neighbourhood which will help us use local observations to answer some of the key questions about the young universe."
    },
    "1005/1005.4375_arXiv.txt": {
        "abstract": "{The \\ion{Na}{x} X-ray lines between 10.9 and 11.2~\\AA\\ have attracted little attention but are of interest since they enable an estimate of the coronal abundance of Na to be made. This is of great interest in the continuing debate on the nature of the FIP (first ionization potential) effect. } {Observations of the \\ion{Na}{x} lines with the Solar Maximum Mission Flat Crystal Spectrometer and a rocket-borne X-ray spectrometer are used to measure the Na/Ne abundance ratio, i.e. the ratio of an element with very low FIP to one with high FIP.} {New atomic data are used to generate synthetic spectra which are compared with the observations, with temperature and the Na/Ne abundance ratio as free parameters.} {Temperature estimates from the observations indicate that the line emission is principally from non-flaring active regions, and that the Na/Ne abundance ratio is $0.07 \\pm 50$\\%.} {The Na/Ne abundance ratio is close to a coronal value for which the abundances of low-FIP elements (FIP $< 10$~eV) are enhanced by a factor of 3 to 4 over those found in the photosphere. For low-temperature ($T_e \\leqslant 1.5$~MK) spectra, the presence of \\ion{Fe}{xvii} lines requires that either a higher-temperature component is present or a revision of ionization or recombination rates is needed. } ",
        "introduction": "Sodium has few prominent or unblended emission lines in the ultraviolet and X-ray spectra emitted by solar coronal plasmas. Hence determinations of its abundance relative to other elements rely on measurements of emission line fluxes that are either weak or in crowded spectral regions. Yet its abundance is important for discussions of the so-called FIP effect, whereby coronal element abundances differ from corresponding photospheric values depending on whether their first ionization potential (FIP) is greater or less than about 10~eV \\citep{mey85,fel92a}. This is because the FIP of Na is only 5.14~eV, less than any other element common in the Sun apart from potassium, which has an FIP of 4.34~eV. Determinations of the coronal abundance of either Na or K can therefore probe the possibility that elements with small FIP are enhanced by amounts that depend on the magnitude of the FIP, as is suggested by the electric current model of \\citet{hen97} to explain the FIP effect. Observational evidence for this has been provided by \\citet{fel93}, and to some extent by observations of He-like K X-ray lines seen in numerous flares and non-flare periods between 2002 and 2005 with the RESIK spectrometer \\citep{syl10}. Coronal sodium emission lines include the Li-like Na (\\ion{Na}{ix}) resonance lines $1s^2 2s\\,\\,^2S_{1/2}-1s^2\\,2p\\,\\,^2P_{1/2}$, $^2P_{3/2}$ in the extreme ultraviolet, at 681.7~\\AA\\ and 694.3~\\AA, and the He-like Na (\\ion{Na}{x}) line $1s2s\\,\\,^3S_1 - 1s2p\\,\\,^3P_2$ at 1111.76~\\AA\\ seen in {\\it SOHO} SUMER spectra \\citep{cur00, cur01}, and resonance lines of H-like (\\ion{Na}{xi}) and He-like Na (\\ion{Na}{x}) in the soft X-ray range. X-ray observations of the \\ion{Na}{x} and \\ion{Na}{xi} X-ray emission lines should in principle be easier than for the equivalent potassium lines as the photospheric abundance of Na is higher by a factor of 16 than K \\citep{asp09}, but nevertheless observations remain scarce. The \\ion{Na}{xi} Ly-$\\alpha$ and Ly-$\\beta$ lines are at 10.02~\\AA\\ and 8.46~\\AA\\ respectively, and the \\ion{Na}{x} $1s^2 - 1s2l$ ($l=s$, $p$) lines are near 11~\\AA. The Ly-$\\beta$ line was observed by \\citet{wal74} with crystal spectrometers on the {\\it OV1-17} spacecraft, and the \\ion{Na}{x} lines with rocket-borne crystal spectrometers by \\citet{par75}. Only a single scan of the Flat Crystal Spectrometer (FCS) on {\\it Solar Maximum Mission} ({\\it SMM}) in its 9-year lifetime was made of the \\ion{Na}{x} lines during a flare in 1980 \\citep{phi82}. Previous Na abundance analyses of the {\\it OV1-17} data gave ${\\rm log}\\,A({\\rm Na}) = 6.26$ \\citep{wal74} and of the Parkinson rocket data ${\\rm log}\\,A({\\rm Na}) = 6.73$ \\citep{par75} (on a scale ${\\rm log}\\,A({\\rm H}) = 12)$. The \\ion{Na}{x} and \\ion{Na}{xi} X-ray lines have incidentally been observed in radiation from a laboratory (Z-pinch) device by \\citet{bur90}.  Only rather approximate atomic data were available at the time of the \\citet{wal74} and \\citet{par75} analyses. However, significant advances have since been made to evaluate collisional excitation rate coefficients and ionization fractions needed for the evaluation of the abundance of sodium from X-ray line fluxes. The {\\sc chianti} database and software package in the Interactive Data Language (IDL) SolarSoftWare system, since its inception in the late 1990s \\citep{der97, der09}, has also simplified analyses of X-ray and ultraviolet spectral data considerably. Excitation data for the \\ion{Na}{x} lines presently included in {\\sc chianti} are interpolated values from other elements, but more recently \\citet{agg09} have calculated collisional excitation data specifically for He-like sodium with the close-coupling $R$-matrix code. Also, more refined ionization fractions assuming coronal ionization equilibrium have recently been calculated by \\citet{bry09}.  The availability of new atomic data has inspired a fresh analysis, reported here, of the \\ion{Na}{x} X-ray lines seen in the {\\it SMM} FCS spectral scan and the spectra obtained by \\citet{par75}. X-ray lines of \\ion{Na}{x} were identified by \\citet{phi82}, with more detailed analysis of nearby lines by \\citet{lan05}. A slight revision of spectral line identifications is given here and a determination of the Na/Ne abundance ratio for this flare plasma, based on a remarkable coincidence of the contribution functions of the \\ion{Na}{x} $w$ line and that of the \\ion{Ne}{ix} $1s^2-1s4p$ ($w4$ or He-$\\gamma$) line blended with the \\ion{Na}{x} $w$ line. The Na/Ne abundance ratio is examined in the light of other recent abundance determinations and the nature of the FIP effect in solar coronal plasmas. ",
        "conclusions": "In this work, recent collisional excitation results for \\ion{Na}{x} X-ray line emission have been applied to solar spectra for the first time, together with new atomic data for \\ion{Na}{ix} satellites, formed by both dielectronic recombination and inner-shell excitation. From these and other data from the {\\sc chianti} atomic code, synthetic spectra were calculated as a function of electron temperature $T_e$. To date, the only observed spectra taken with high-resolution spectrometers in the 10.9--11.3~\\AA\\ range are from the {\\it SMM} Flat Crystal Spectrometer (a single scan during the decay of a M1.5 flare) and two rocket-borne scanning crystal spectrometers in 1971 viewing a non-flaring active region. Comparison of the FCS and calculated spectra in the 10.9--11.3~\\AA\\ range shows that the \\ion{Na}{x} X-ray lines with nearby \\ion{Na}{ix} dielectronic satellites yield a Na/Ne abundance ratio through the blend of the \\ion{Na}{x} $w$ (resonance) line with the \\ion{Ne}{ix} $w4$ ($1s^2\\,\\,^1S_0 - 1s4p\\,\\,^1P_1$) line, these lines having almost identical contribution functions $G(T_e)$ for the low temperatures considered here. The value obtained, $A({\\rm Na})/A({\\rm Ne}) = 0.07 \\pm 50$\\%, is the coronal value of \\citet{fel92a} to within uncertainties. It is significantly higher than the photospheric value, 0.020, and suggests that sodium (like other low-FIP elements) is enhanced by a factor of 3 to 4 in the corona. The lower-temperature observations obtained by \\citet{par75} are best fitted with theoretical spectra having $T_e = 1$~MK, and are again consistent with a coronal Na/Ne abundance ratio. The presence of \\ion{Fe}{xvii} lines in these spectra is unexpected, the most recent ionization equilibrium calculations of \\citet{bry09} indicating either that there is a higher-temperature component present or (more likely) that there is a negligible fraction of Fe$^{+16}$ ions at $T_e \\leqslant 1$~MK. A revision of ionization or recombination rates may be needed."
    },
    "1005/1005.3249_arXiv.txt": {
        "abstract": "{The fundamental properties   of brown dwarfs evolve  with age. Models describing  the evolution of  luminosities and effective temperatures, among other physical parameters, can  be empirically constrained using brown dwarfs of various masses in star clusters of well determined age  and metallicity.} {We aim to carry out a spectroscopic and photometric characterization of low-mass brown dwarfs  of the $\\sim$120~Myr old \\object{Pleiades} open cluster.} {We obtained low-resolution near-infrared spectra of the $J$ = 17.4--18.8~mag candidate L-type brown dwarfs \\object{PLIZ~28} and 35, \\object{BRB~17}, 21, 23, and 29, which are Pleiades members by photometry and proper motion. We also obtained spectra of the well-known $J$ = 15.4--16.1~mag  late M-type cluster members \\object{PPl~1}, \\object{Teide~1}, and \\object{Calar~3}.} {We find that the former six objects have early- to mid-L spectral types and confirm previously reported M-types for the other three objects. The spectra of the L0-type BRB~17 and PLIZ~28 present a triangular $H$-band continuum shape, indicating that this peculiar spectral feature persists until at least the age of the Pleiades. We add to our sample 36 reported M5--L0-type cluster members, collecting their $I_{\\rm C}$- and UKIDSS $ZYJHK$-band photometry. We confirm a possible interleaving of the Pleiades and field L-type sequences in the $JHK$ absolute magnitude versus spectral type diagrams, and quantify marginally redder Pleiades $J-K$ colours, by 0.11$\\pm$0.20~mag, possibly related to both reddening and youth. Using field dwarf bolometric correction -- and effective temperature -- spectral type relations, we obtain the Hertzsprung--Russell diagram of the Pleiades sample.  Theoretical models reproduce well the spectral sequence at M5.5--9, but appear to overestimate the luminosity or underestimate the effective temperature at L0--5.} {We classify six faint Pleiades brown dwarfs as early to mid L-type objects using low-resolution near-infrared spectra. We compare their properties to field dwarfs and theoretical models and estimate their masses to be in the range 0.025--0.035~$M_{\\odot}$.} ",
        "introduction": "About 600 field L-type dwarfs have been spectroscopically identified to date.\\footnote{http://spider.ipac.caltech.edu/staff/davy/ARCHIVE/index.shtml,\\\\ 2010 May 17.} Most of them have been discovered by large sky-area surveys such as 2MASS \\citep{skrutskie2006} and SDSS \\citep{abazajian2009}. Spectroscopy has allowed the determination of physical properties and the derivation of effective temperatures (in the range $\\sim$2300--1400~K; e.g.\\ \\citealt{leggett2001}), and parallax studies have provided distances and luminosities for about 10\\% of these objects, indicating that they are most probably very low mass stars or massive brown dwarfs. For individual field L-dwarfs, an estimation of mass from the luminosity and effective temperature requires additional information on age, which is in general difficult to obtain. This limits our ability to determine the field substellar mass function and formation history \\citep{reid1999,chabrier2002,chabrier2003,kroupa2003,burgasser2004c,allen2005,pinfield2006}.    \\begin{table*} \\caption{Spectroscopic observations.} \\label{log} \\centering \\begin{tabular}{c c c c c c c c c c} \\hline\\hline Object      & $J$\\tablefootmark{b}\t      & Opt.&  Instrument\\tablefootmark{d}   & Date\t& $T_{\\rm exp}$\\tablefootmark{e} & Airmass & $\\Delta$airmass$_{\\rm T-S}$ & S/N\\tablefootmark{f} & Near-IR\t    \\\\ &  (mag)\t                      & SpT\\tablefootmark{c}\t\t    &\t\t                 &\t& (min) \t                 &\t   &\t\t                 &                      & SpT\\\\ \\hline \\object{PPl~1}       & 15.36$\\pm$0.01  & M6.5    \t&  NICS($JHK$)    & 2002 Oct 31 & 8$\\times$2         & 1.08       & 0.14 & 50\t\t & M7.0$\\pm$0.5  \\\\ \\object{Calar~3}     & 16.08$\\pm$0.01  & M8      \t&  NICS($JHK$)    & 2002 Oct 31 & 8$\\times$2         & 1.15       & 0.07 & 49\t\t & M8.0$\\pm$0.5  \\\\ \\object{Teide~1}     & 16.21$\\pm$0.01  & M8      \t&  NICS($JHK$)    & 2002 Nov 27 & 8$\\times$2         & 1.03       & 0.73 & 32\t\t & M8.0$\\pm$0.5  \\\\ \\object{BRB~17}      & 17.41$\\pm$0.03  & ...\t\t&  LIRIS($zJ,HK$) & 2006 Dec 28 & (12, 18)$\\times$2  & 1.74, 1.41 & 0.44, 0.89 & 4, 6\t & L0.0$\\pm$1.0 \\\\ \\object{PLIZ~28}     & 17.60$\\pm$0.04  & ...\t\t&  NICS($JHK$)    & 2005 Nov 20 & 15$\\times$2\t     & 1.11       & 0.57 & 18\t\t & L0.0$\\pm$1.0 \\\\ \\object{PLIZ~35}     & 18.07$\\pm$0.05  & ...\t\t&  NICS($JHK$)    & 2005 Sep 18 & 65$\\times$0.6      & 1.03       & $-$0.03 & 15\t & L2.0$\\pm$1.0 \\\\ \\object{BRB~21}      & 18.14$\\pm$0.06  & ...\t\t&  NICS($JHK$)    & 2005 Sep 18 & 80$\\times$0.5\t     & 1.21       & 0.15 & 7\t\t & L3.0$\\pm$1.0 \\\\ \\object{BRB~23}      & 18.23$\\pm$0.04  & ...\t\t&  NICS($JHK$)    & 2007 Feb 7  & 10$\\times$5\t     & 1.33       & $-$0.27 & 14\t & L3.5$\\pm$1.0 \\\\ \\object{BRB~29}      & 18.69$\\pm$0.10  & ...\t\t&  LIRIS($HK$)    & 2006 Dec 29 & 12$\\times$2\t     & 1.09       & $-$0.02 & 3\t\t & L4.5$\\pm$1.5 \\\\ \\hline J0829\\tablefootmark{a} & 12.80$\\pm$0.03  & L2V     \t&  LIRIS($zJ,HK$) & 2006 Dec 28 &(4, 6)$\\times$1.5   & 1.38, 1.36 & 0.80, 0.94 & 29, 37& L2.0$\\pm$1.0  \\\\ \\hline \\end{tabular} \\begin{flushleft} \\tablefoottext{a}{\\object{SSSPM~J0829$-$1309}.} \\tablefoottext{b}{From the Galactic Clusters Survey (GCS) component of UKIDSS \\citep[][sixth data release]{lawrence2007} for the Pleiades members and from 2MASS \\citep{skrutskie2006} for the field object.} \\tablefoottext{c}{From \\citet{martin1996} for the three Pleiades objects and from \\citet{scholz2002} and \\citet{lodieu2005} for the field object.} \\tablefoottext{d}{A~slit width of 1.0~arcsec was used, except for Calar~3 and PPl~1 (1.5~arcsec). The NICS/Amici spectra have a constant resolving power $R=\\lambda/\\Delta\\lambda\\approx50$ in the range 0.8--2.5~$\\mu$m (dispersion 3--10~nm~pix$^{-1}$); for Calar~3 and PPl~1, $R\\approx33$. The LIRIS $zJ$-band (0.887--1.531~$\\mu$m) spectra have $R\\approx960$ at $\\lambda\\approx1.17$~$\\mu$m and a dispersion of 0.61~nm~pix$^{-1}$, whereas the $HK$-band (1.388--2.419~$\\mu$m) spectra have $R\\approx945$ at $\\lambda\\approx1.83$~$\\mu$m and a dispersion of 0.97~nm~pix$^{-1}$.} \\tablefoottext{e}{Dithers $\\times$ detector integration time.} \\tablefoottext{f}{Pseudo-continuum signal-to-noise ratios are measured in the interval 2.14--2.24~$\\mu$m, or for the LIRIS zJ-band data, in the interval 1.28--1.32~$\\mu$m.} \\end{flushleft} \\end{table*}     An empirical determination of the evolution of effective temperature and luminosity with age for L-type dwarfs could be obtained by identifying such objects in stellar associations of various ages. Brown dwarf- and planetary-mass candidates of (mostly early) L-type have been studied photometrically and spectroscopically in the \\object{Serpens} \\citep{lodieu2002}, \\object{Ophiuchus} \\citep{jayawardhana2006}, \\object{Chamaeleon~I} \\citep{luhman2006,luhman2008b,luhman2008c,schmidt2008}, \\object{Chamaeleon~II} \\citep{jayawardhana2006b,allers2007}, \\object{Taurus} \\citep{itoh2005,luhman2009}, \\object{Trapezium} \\citep{lucas2001,lucas2006,weights2009}, and \\object{Lupus~1} clouds \\citep{neuhauser2005,mcelwain2007}, the \\object{$\\sigma$~Orionis} open cluster \\citep{zapateroosorio1999Or,zapateroosorio2000,bejar2001,martin2001,barradoynavascues2001b,barradoynavascues2002or,mcgovern2004}, and the \\object{TW~Hydrae} \\citep{gizis2002,chauvin2004,looper2007}, \\object{Upper~Scorpius} \\citep{lodieu2008,lafreniere2008,bejar2008}, and \\object{Tucana--Horologium} associations \\citep{chauvin2005b}. All these star-forming regions are very young, with ages of a few Myr to a few tens of Myr. Finally, several L-type dwarfs characterized photometrically and spectroscopically have been associated with moving groups, such as the all-sky, 400$\\pm$100~Myr old Ursa Major moving group \\citep{bannister2007,jameson2008a,casewell2008}.  In this context, the \\object{Pleiades} open cluster is well suited to provide a reference sequence for L- and T-dwarfs at an age of $\\sim$120 Myr. Proper motion studies \\citep{moraux2003,bihain2006,lodieu2007} have shown that there is a significant population of low-luminosity members of the cluster that, based on their optical and near-infrared colours, are likely to be of  spectral type L. \\citet{casewell2007} have extended the brown dwarf search beyond these objects with some candidate members of  spectral type T. Here, we report near-infrared spectroscopy for six of the nine L-type candidates confirmed as Pleiades proper motion members by \\citet{bihain2006}, which allows us to study the relationship between spectral type and luminosity for L dwarfs at a homogeneous age and solar metallicity. ",
        "conclusions": "We obtained near-infrared NICS/Amici and LIRIS spectra of six $J=17.4$--18.8~mag candidate L-type brown dwarfs of the $\\sim$120~Myr old Pleiades cluster, as well as Amici spectra of three other members that are well-known late M-type brown dwarfs. We confirm that the former objects have early- to mid-L spectral types and also confirm the optical spectral types reported for the latter objects. \\object{PLIZ~28} and \\object{BRB~17}, the two earliest L-type brown dwarfs, present a triangular $H$-band continuum shape and a noticeable VO band, associable to low surface gravity, dust and youth. The Pleiades L-type brown dwarfs appear to have similar absolute magnitudes and colours as their field counterparts, but have slightly redder near-infrared colours, possibly related to both reddening and youth.   We build the Hertzsprung--Russell diagram of our Pleiades sample using field dwarf relations and find good agreement with theoretical models for the M5.5--9~type objects. However, for the L0--5~type brown dwarfs, models  appear to overpredict the luminosity ($\\Delta \\log{L/L_{\\odot}}=0.1-0.4$~dex) or underestimate the effective temperature ($\\Delta T_{\\rm eff}=100$--400~K). It is also possible that  the adopted field dwarf relations are not truly valid for $\\sim$120~Myr-old objects. Additional optical-to-infrared spectroscopy and higher resolution imaging of the Pleiades L-type brown dwarfs may allow us to account more for dust opacities and binarity \\citep[e.g.,][]{knapp2004,burgasser2008}, and refine our estimates of spectral types, effective temperatures, luminosities, and masses. Current models predict masses in the range 0.025--0.035~$M_{\\odot}$ for the Pleiades L0--5~type brown dwarfs."
    },
    "1005/1005.1882_arXiv.txt": {
        "abstract": " ",
        "introduction": "It is generally believed that QCD exhibits a rich phase structure in an extreme environment such as at high temperature and high baryon chemical potential. For the chiral phase transition, it is a widely accepted view that the critical point(s) (i.e.\\ the end point of the first-order phase boundary) should exist at finite temperature and density. Usually, a schematic $T$-$\\mu$ phase diagram with one critical point is widely adopted in the literature~\\cite{Stephanov:2007fk}.  The color-superconducting (CSC) phase may appear in the $T$-$\\mu$ phase diagram of QCD for low temperature and large baryon chemical potential~\\cite{WilczekReview,RischkeReview,BuballaReview,AlfordReview}. At asymptotically high density that justifies the perturbative QCD calculations, the color-flavor locked (CFL) phase~\\cite{CFL} has been established as the ground state of quark matter. For the intermediate density region which may exist in the core of a compact star, the nonperturbative features of QCD play a more important role in the phase structure of QCD, and both CFL and non-CFL CSC phases may manifest themselves in this region.  The appearance of the CSC phase around the boundary of the chiral phase transition at low temperatures should affect the chiral phase transition,  which may result in an unexpected phase structure of QCD.  Especially, the competition between the chiral condensate and the diquark condensate may lead to the emergence of the new chiral critical point(s) in the low temperature region. Such an example was first presented in~\\cite{KitazawaVector} in a two-flavor Nambu-Jona-Lasinio (NJL) model with the vector interaction : It was found that the repulsive vector interaction can lead to a two-critical-point structure in the $T$-$\\mu$ phase diagram of QCD.  The appearance of the new critical point is attributed to the fact that the density-density correlation induced by the vector interaction effectively enhances the competition mentioned above while weakening the first-order chiral restoration. For the three flavor case, the realization of a similar QCD phase diagram has been recently conjectured ~\\cite{Hatsuda:2006ps}, where the $\\mathrm{U_A}(1)$-breaking vertex may induce a new critical point in the low temperature region. This conjecture is based on a general Ginzburg-Landau theory constrained by QCD symmetries and the CFL phase is assumed to appear near the chiral boundary in the $SU_f(3)$ limit. It has been argued that the resultant crossover of the chiral restoration at small temperatures embodies the hadron-quark continuity hypothesis~\\cite{Continuity}.   Note that the two-critical point structures obtained above are all based on the  exact flavor symmetry. Namely, the quarks with different flavor have the same mass and density. In realistic situations, $s$ quark is much heavier than $u$ and $d$ quarks, which may make a different phase structure involving the chiral transition. In particular, the large mass difference between $s$ quark and $u(d)$ quark may disfavor the formation of the CFL in the intermediate density region. Thus, the conjecture of the two-critical-point structure claimed in ~\\cite{Hatsuda:2006ps} may not be true in the realistic situation. In addition,  near the chiral boundary, even though the mass difference between the $u$ quark and the $d$ quark is very small in contrast to the quark chemical potential,  the constraint by local electric-charge neutrality may lead to a relatively large density disparity between them. In short,  we can expect that the Fermi sphere differences among $u$, $d$ and $s$ quarks should have significant influence on the chiral phase transition in the presence of CSC.  Considering the local charge neutrality constraint, it was first disclosed in ~\\cite{Zhang:2008wx} that the positive electric chemical potential $\\mu_e$ plays a similar role on the chiral transition as the repulsive vector interaction. This is another mechanism which can realize a multiple chiral critical-point structure.  In a simple two-flavor NJL model, the same two-critical-point structure found in ~\\cite{KitazawaVector} can be realized without introducing vector interaction when taking into account the local electric-neutrality~\\cite{Zhang:2008wx}. Moreover, besides its role as an effective vector interaction, the positive $\\mu_e$ also implies a finite disparity between Fermi spheres of $u$ and $d$ quarks. For a system composed of two different Fermions with the mismatched Fermi spheres, it is known that the energy gap of the Cooper pairing between them can {\\em increase} with temperature. This unusual thermal behavior is due to the smearing of the Fermi surface by the temperature. Such an unconventional behavior in the two-flavor neutral CSC phase can lead to a special competition between the chiral condensate and the diquark condensate, which may be enhanced with increasing temperature. For some model parameter regions, this abnormal competition can give rise to {\\em three} chiral critical structure in the phase diagram~\\cite{Zhang:2008wx}. We remark that the interplay of two order parameters with the external constraint(s) should be a general mechanism for a realization of a multiple critical-point structure in many-body system including those in condensed matter physics.   It is an interesting problem whether the multiple critical-point structure can appear in the QCD phase diagram for the two-plus-one-flavor case. The CFL phase may be disfavored near the chiral boundary in this case due to the relatively heavy strange quark. In addition, the $\\mathrm{U_A}(1)$ anomaly of QCD may have significant influence on the chiral phase transition. Such an anomaly can be successfully fulfilled in the effective quark field theory with the so-called Kobayashi-Maskawa-'t Hooft(KMT) term\\cite{KMT}. Because of the similar roles of the repulsive vector interaction and the electric chemical potential under the neutrality constraint on the chiral phase transition, it is naturally expected that chiral restoration is weakened more significantly when taking into account both aspects simultaneously. The study of the QCD phase diagram including all these ingredients has been done in \\cite{Zhang:2009zk} in the framework of a two-plus-one-flavor NJL model.   As we will show below, the multiple chiral critical point structures can still emerge in the QCD phase diagram for the two-plus-one-flavor case. Such a conclusion is qualitatively different from the result in \\cite{Ruester:2005jc,Abuki:2005ms} where the same model was used without considering the vector interaction. This just reflects the combined effect of the electric chemical potential and vector interaction on the chiral phase transition through a special competition of the chiral condensate and the diquark condensate. It should be stressed that the present mechanism giving rise to the multiple chiral critical structures is different from that reported in \\cite{Hatsuda:2006ps} where the anomaly term in the flavor-$SU(3)$ limit plays the essential role. In our case,  only two-flavor CSC (2CSC) phase appears near the chiral boundary and there is no cubic interactions between the chiral condensate and the diquark condensate and the number of the chiral critical points can be more than two. ",
        "conclusions": "The vector-vector interaction and electric-charge neutrality in $\\beta$ equilibrium have important effects on the chiral phase transition. In the presence of color superconductivity(CSC), as demonstrated in \\cite{Zhang:2009zk} within the NJL model,  the QCD phase diagram may have a multiple chiral critical-point structure. The emergence of such a multiple critical-point structure is made robust with the inclusion  of the repulsive vector coupling $G_V$ between quarks, and it becomes that there always exists a parameter window in the NJL models which favors the appearance of a multiple chiral critical-point structure for a wide range of the vacuum quark mass, i.e., from 300 MeV to 400 MeV \\cite{Zhang:2009zk}. More precisely, besides the two- and three-critical-points structures found in \\cite{KitazawaVector,Zhang:2008wx} and  \\cite{Zhang:2008wx}, respectively, the QCD phase diagram can have even {\\em four} critical points; such a multiple critical-point structure is caused by the combined effect of the positive $\\mu_e$ and $G_V$. Because the dynamical strange quark mass is still relatively large near the boundary of the chiral transition, the multiple critical-point structures found in the two-flavor case also appear in the two-plus-one-flavor case. For the intermediate diquark coupling,  the number of critical points changes  as $1\\, \\rightarrow \\, 2\\,\\rightarrow\\, 4\\,\\rightarrow\\, 2\\, \\rightarrow\\,0$ with an increasing vector coupling in the two-plus-one-flavor NJL model. In general, one can expect that different model parameters may possibly give other order of the number of the critical points as the vector coupling is increased.  Although our analysis is based on a low-energy effective  model which inherently has, more or less,  a parameter dependence, we have seen that the physical mechanism to realize the multiple critical-point structure is solely dependent on the basic ingredients of the effective quark dynamics and thermodynamics. Therefore, we believe that the results obtained in the present work should be taken seriously and examined in other effective models of QCD,  or hopefully lattice QCD simulations.  Our result also has a meaningful implication for the study of phase transitions in condensed matter physics. That means some external constraints enforced on the system can lead to the formation or expansion of the coexisting phase,  and the competition between two order parameters can give rise to multiple critical points.  Last but not least, in \\cite{Zhang:2009zk} we have shown for the first time that the repulsive vector interaction  does suppress the chromomagnetic instability related to the asymmetric homogeneous 2CSC phase. Even such a conclusion is based on a nonlocal two-flavor NJL model, it should also be true for the two-plus-one-flavor case because of the relatively large mass of s quark. With increasing vector interaction, the unstable region associated with chromomagnetic instability shrinks towards lower temperatures and higher chemical potentials. This means that the vector interaction can partially or even totally resolve  the chromomagnetic instability problem.  To cure the chromomagnetic instability, inhomogeneous asymmetric color superconductivity phases such as the LOFF phase and the gluonic phase were proposed in the literature. For the inhomogeneous phase, beside the condensate $\\langle\\overline{\\psi}{\\gamma_0}\\psi\\rangle$, there is no reason to rule out the appearance of another new condensate , $\\langle\\overline{\\psi}\\vec{\\gamma}\\psi\\rangle$,  when considering the vector interaction. The effect of both the timelike vector condensate and the spacelike condensate on the asymmetric inhomogeneous  CSC phase will be reported in our future work.  \\bigskip"
    },
    "1005/1005.4682_arXiv.txt": {
        "abstract": "This paper presents late-time near-infrared and {\\it Spitzer} mid-infrared photometric and spectroscopic observations of warm dust in the Type IIn SN 2005ip in NGC 2906.  The spectra show evidence for two dust components with different temperatures.  Spanning the peak of the thermal emission, these observations provide strong constraints on the dust mass, temperature, and luminosity, which serve as critical diagnostics for disentangling the origin and heating mechanism of each component.  The results suggest the warmer dust has a mass of $\\sim 5 \\times 10^{-4}~$\\msolar, originates from newly formed dust in the ejecta, or possibly the cool, dense shell, and is continuously heated by the circumstellar interaction.  By contrast, the cooler component likely originates from a circumstellar shock echo that forms from the heating of a large, pre-existing dust shell $\\sim 0.01 - 0.05$~\\msolar~by the late-time circumstellar interaction.  The progenitor wind velocity derived from the blue edge of the He 1 1.083 \\micron~P Cygni profile indicates a progenitor eruption likely formed this dust shell $\\sim$100 years prior to the supernova explosion, which is consistent with a Luminous Blue Variable (LBV) progenitor star. \\\\ \\\\ \\\\ ",
        "introduction": "\\label{sec3:intro}  SN 2005ip was discovered in NGC 2906 ($d \\approx$ 30 Mpc) on UT 2005 November 5 \\citep{boles05}.  Early optical spectra suggested the discovery occurred a few weeks following the explosion \\citep{modjaz05}.  The supernova is a Type IIn event given the development of narrow H$\\alpha$~emission lines \\citep{smith09ip}.  \\citet{fox09} discovered a late-time near-infrared excess that has persisted at least two and a half years post-discovery.  Late-time infrared emission typically indicates the presence of warm dust.  The origin and heating mechanism of the dust, however, is not always well constrained.  The dust may be newly formed or may have existed at the time of the supernova.  If newly formed, the dust may condense from refractory elements in the expanding supernova ejecta or in the cool, dense shell of post-shocked  circumstellar gas lying in between the forward and reverse shocks (like the post-shocked environment trailing stellar wind collisions in WR binary systems \\citep{usov91}).  In both cases, several heating mechanisms are possible, including radioactivity, optical emission from circumstellar interaction, and collisional heating by hot gas in the reverse shock.  Alternatively, pre-existing dust may be collisionally heated by hot, shocked gas or radiatively heated by either the peak supernova luminosity or late-time optical emission from circumstellar interaction.  In the latter case, an `IR echo' is evident due to light travel time effects.  Multiple scenarios can contribute to the late-time infrared flux, as in the cases of SNe 2004et \\citep[][Sugerman et al. in prep]{kotak09} and 2006jc \\citep{mattila08}.  Type IIn supernovae, named for their ``narrow'' emission lines \\citep{schlegel90}, are more often observed to exhibit late time infrared emission associated with warm dust than any other supernova subclass \\citep[e.g.,][]{pastorello02,gerardy02, pozzo04, fox09, smith09gy, miller10gy, miller10}.  Disentangling the origin and heating mechanism of the dust can yield important diagnostics concerning the circumstellar medium, supernova progenitor, and explosion dynamics.  For example, if the dust is shock heated, the observed dust temperature yields the gas density \\citep{dwek87,dwek08}, which can be used to trace the progenitor's mass loss history \\citep[e.g.,][]{smith09ip}.  If the dust is newly formed, the observed dust mass can be compared to models that predict supernovae as primary sources of dust at high redshifts \\citep{todini01,nozawa03,nozawa08}.  \\begin{deluxetable*}{c c c c c c c c} \\tablewidth{0pt} \\tablecaption{{\\it Spitzer} Observations Summary \\label{tab1}} \\tablecolumns{8} \\tablehead{ \\colhead{Target} & \\colhead{Position} & \\colhead{Flux} & \\colhead{} & \\colhead{Int/} & \\colhead{Total} & \\colhead{Ramp $\\times$} & \\colhead{\\# of}\\\\ \\colhead{Field} & \\colhead{RA} & \\colhead{Density} & \\colhead{Bands} & \\colhead{Pixel} & \\colhead{Duration} & \\colhead{Cycles} & \\colhead{AORS}\\\\ \\colhead{} & \\colhead{DEC} & \\colhead{} & \\colhead{} & \\colhead{(secs)} & \\colhead{(secs)} & \\colhead{(sec $\\times$ \\#)} & \\colhead{} } \\startdata SN 2005ip & 09:32:06 & 0.4 mJy & IRAC & 300 & 1062 & - &  1 \\\\ & +08:26:44& (8\\micron)& all& & &  &\\\\ & & & & &  & &\\\\ SN 2005ip &  09:32:06 &  0.4 mJy & IRS & - & 9516 & 60 $\\times$ 12 &  1 \\\\ &  +08:26:44& (8\\micron) & Spectral Mapping & & & & \\\\ & & & low short both & & &  &\\\\ \\enddata \\end{deluxetable*}  The origin of the observed dust is not always obvious.  For SN 2005ip, optical spectra show a progressive attenuation of the red wing of both the broad and intermediate lines, directly confirming the formation of new dust in both the ejecta and post-shocked cool, dense shell \\citep{smith09ip}.  Such direct evidence, however, is rare.  No more than a handful of supernovae \\citep{lucy91,meikle93,elmhamdi04,elmhamdi03,sugerman06,kotak08,pozzo04,smith09ip,smith08jc} show direct evidence of dust formation in the ejecta, and aside from SN 2005ip, only SNe 1998S \\citep{pozzo04} and 2006jc \\citep{smith08jc} show direct evidence of dust formation in the cool, dense shell (although this scenario is invoked to explain the asymmetries observed in the spectra of SNe 2004et \\citep{kotak09} and 2007od  \\citep{andrews10}).  The observed dust yields, however, all tend to be 2-3 orders of magnitude lower than required to account for the large amounts of dust observed at high redshifts.  Mid-infrared ($3 \\le \\lambda \\le 15$~\\micron) observations span the peak of the thermal spectral energy distribution from dust with temperatures ranging $100 \\lesssim T_d \\lesssim 1000$~K, providing strong constraints on the dust mass, temperature, and, thereby, the luminosity.  These quantities serve as useful diagnostics for disentangling the origin and heating mechanism of warm dust.  Late-time mid-infrared observations of supernovae, particularly the Type IIn subclass, are rare.  In this paper, we present the first late-time (day 936 post-discovery) $Spitzer$/IRS spectra of a Type IIn supernova, as well as coincident $Spitzer$/IRAC photometry.  We also present 0.9-2.5 \\micron, $R=3000$~spectra obtained with APO 3.5-m/TripleSpec.  Section \\ref{sec3:obs} presents the observations and data reduction techniques. The combined spectra show evidence for two independent dust components: a hot, near-infrared (HNI) and warm, mid-infrared (WMI) component.  For each component, we derive the dust composition, mass, temperature, and luminosity.  Section \\ref{sec3:source} explores the origin and heating mechanism of these components to determine the degree to which SN 2005ip forms new dust.  The HNI dust mass originates predominantly from newly formed dust in the ejecta, while the WMI component likely originates from an `IR echo' in a pre-existing dust shell.  We use these results to explore the progenitor system and its evolution.  Section \\ref{sec3:con} presents a summary of the findings and a discussion of future work. ",
        "conclusions": "\\label{sec3:con}  The $Spitzer$~spectrum presented here confirms the presence of warm dust in SN 2005ip.  Combined with near-infrared observations, the results show evidence for two independent dust masses: a hot, near-infrared (HNI) and warm, mid-infrared (WMI) component.  Infrared observations span the peak of the thermal emission, thereby providing strong constraints on the dust mass, temperature, and luminosity, which serve as critical diagnostics for disentangling the origin and heating mechanism of each component.  The HNI dust mass originates primarily from newly formed dust in the ejecta, or possibly the cool, dense shell, while the WMI component likely originates from an circumstellar shock echo that forms from the heating of a large, pre-existing dust shell.  For both components, the heating mechanism is likely the optical luminosity generated from the forward shock interaction with the circumstellar medium.  For wind speeds of $\\sim 100 - 200$~\\kms, the WMI dust shell likely formed via an eruption $\\sim100$ years before the supernova.  These characteristics are consistent with a LBV progenitor, which has been linked to some core-collapse supernovae \\citep{kotakvink06}, but is an emerging trend particularly within the Type IIn subclass \\citep[e.g.,][]{gal-yam07,smith07, trundle08, smith08tf, smith09gy, gal-yam09}.  The newly formed HNI dust mass, $M_d (\\rm HNI) \\sim 5 \\times 10^{-4}~$\\msolar, is at least two orders of magnitude lower than predicted by ejecta condensation models to reproduce the large amounts of dust observed at high redshifts \\citep{todini01,nozawa03,nozawa08}.  A dust mass of $\\sim 10^{-4}~$\\msolar~is comparable to other Type IIn events \\citep[e.g.,][]{gerardy02, pozzo04, meikle07}.  For clumpy ejecta, it should be noted that both semi-analytical models \\citep{varosi99} and Monte Carlo radiative-transfer simulations \\citep{ercolano07} have shown that dust masses can be at least an order of magnitude larger than that predicted by the analysis presented in Section \\ref{sec3:dust}.  The clump structure for any supernova, however, remains unconstrained at present.  The pre-existing WMI dust mass, $M_d (\\rm WMI) \\sim 0.05$~\\msolar, is quite significant.  Although, in the case of SN 2005ip, the forward shock may ultimately destroy this dust (mid-infrared observations at later epochs will reveal the accuracy of this model), the forward shocks may be sufficiently decelerated in other systems to allow for dust survival.  Late-time mid-infrared observations of dust in the supernova environment provide for a unique interpretation of the circumstellar environment and progenitor system.  Nonetheless, mid-infrared observations are quite rare.  Presented in this paper is the first mid-infrared spectrum of any Type IIn supernova.  In the future, we hope to grow the database of mid-infrared observations of Type IIn supernovae.  Doing so, however, can be slow as Type IIn events are particularly rare, consisting of only $\\sim$2-3\\% of all core-collapse supernovae \\citep{gal-yam07} and occurring at a rate of no more than 10/yr out to 150 Mpc \\citep{dahlen99}.  We therefore plan to revisit all Type IIn supernovae from the past ten years to determine the degree to which this subclass exhibits late-time dust emission and identify the emission mechanism.  \\vspace{10 mm}  This work is based on observations made with the Spitzer Space Telescope (PID 50256), which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA. Support for this work was provided by NASA through an award issued by JPL/Caltech.  O. D. F. is grateful for support from NASA GSRP, ARCS, and VSGC.  R. A. C. was supported by NSF grant AST-0807727.  The authors would like to thank Genevieve de Messieres, JD Smith, and the entire Spitzer Science Help Desk for their extensive help with the $Spitzer$~data reduction.  We are also grateful to Joel Sop for his help with the data reduction in the early stages.  We are also thankful to Sarah Doverspike and the anonymous referee for their helpful edits.  \\clearpage"
    },
    "1005/1005.4361_arXiv.txt": {
        "abstract": "s{ The low frequency tail of the CMB spectrum, down along the radio range ($\\sim$1 GHz), may carry weak spectral distortions which are fingerprints of processes occurred during different epochs of the thermal history of the Universe, from $z\\sim 3\\times 10^6$ to reionization. TRIS and ARCADE2 are the most recent experiments dedicated to the exploration of this chapter of CMB cosmology. The level of instrumental accuracy they reached in the determination of the absolute sky temperature is such that the removal of galactic and extra-galactic contamination is the true bottleneck towards the recovery of the cosmological signal. This will be certainly the case also for future experiments in the radio domain. Here we present an update of a study originally done to recognize the contribution of unresolved extra-galactic radio sources to the sky brightness measured by TRIS. Despite the specific context which originated our analysis, this is a study of general interest, improved by the inclusion of all the source counts available up-to-date from 150 MHz to 8.4 GHz.} ",
        "introduction": " ",
        "conclusions": "This update of our previous work allowed us a better reconstruction of differential source counts, especially at faint fluxes. In fact, unlike the first study, we could determine the ratio $A_2/A_1$ finding a hint for weak frequency dependence. Improved source count fits lead directly to a more accurate estimate of the brightness temperature: this is very well described by a single power-law frequency spectrum. Then, using eq.~\\ref{eq:Temp}, we could evaluate the UERS contribution at TRIS frequencies with uncertainties almost a factor 10 smaller than those commonly assumed before new estimates by Gervasi et al.~\\cite{gervasi2}. By virtue of that, the subtraction of this foreground is no longer a limiting factor in low frequency CMB experiments, the overall error budget being now dominated by the reconstruction of the absolute temperature scale and by the disentangling of galactic and cosmological signals. Finally, we conclude that the estimate given by eq.\\ref{eq:Temp}, exclusively based on source counts, doesn't agree with the extragalactic radio excess detected as a diffuse signal by the ARCADE2 collaboration~\\cite{fixsen3}. The discrepancy is around a factor 5 at 600 MHz, and still the origin of the radio excess remains mysterious. If its origin is truly extragalactic, there must be a number of sources able to produce an intense integrated signal, but individually so weak to escape from detection with radio interferometers. The case is intriguing and again tells us how non trivial is searching for CMB spectral distortions, and how unpredictable are the outcomes."
    },
    "1005/1005.1128_arXiv.txt": {
        "abstract": "We observed the ultra-compact binary candidate 4U 0614+091 for a total of 200 ksec with the high-energy transmission gratings onboard the \\chandra X-ray Observatory. The source is found at various intensity levels with spectral variations present. X-ray luminosities vary between 2.0$\\times10^{36}$ \\ergsec and 3.5$\\times10^{36}$ \\ergsec. Continuum variations are present at all times and spectra can be well fit with a powerlaw component, a high kT blackbody component, and a broad line component near oxygen. The spectra require adjustments to the Ne K edge and in some occasions also to the Mg K edge. The Ne K edge appears variable in terms of optical depths and morphology. The edge reveals average blue- and red-shifted values implying Doppler velocities of the order of 3500 \\kms. The data show that Ne K exhibits excess column densities of up to several 10$^{18}$ cm$^{-2}$. The variability proves that the excess is intrinsic to the source. The correponding disk velocities also imply an outer disk radius of the order of $< 10^9$ cm consistent with an ultra-compact binary nature. We also detect a prominent soft emission line complex near the \\oviii L$\\alpha$ position which appears extremely broad and relativistic effects from near the innermost disk have to be included. Gravitationally broadened line fits also provide nearly edge-on angles of inclination between 86 and 89$^{\\circ}$. The emissions appear consistent with an ionized disk with ionization parameters of the order of 10$^4$ at radii of a few 10$^7$ cm. The line wavelengths with respect to \\oviiia\\ are found variably blue-shifted indicating more complex inner disk dynamics. ",
        "introduction": "The low mass X-ray binary (LMXB) \\b0614 has been persistently bright since its detection with \\uhu~\\citep{forman1978}. It has been identified as a type I X-ray burster~\\citep{swank1978} and thus is an accreting neutron star. A recent study of its very bright X-ray bursts lead to a distance estimate of 3.2 kpc~\\citep{kuulkers2009} putting the source a bit beyond earlier projections~\\citep{brandt1992}. It is now confirmed to be an atoll source~\\citep{mendez1997} based on its spectral properties~\\citep{singh1994,vanstraaten2000} and timing pattern ~\\citep{ford1997}. A peculiar trait of its X-ray spectrum is the presence of a hard X-ray tail reaching energies beyond 100 keV \\citep{ford1997} which was modeled as thermal Comptonization at high electron temperatures of greater than 220 keV ~\\citep{piraino1999, fiocchi2008}.  The source has recently become a focus of renewed interest for several reasons. A very recent survey observation covering the source from radio, to infrared, optical, UV and X-ray emissions identified optically thick synchrotron emission from a jet with radiative powers well beyond $10^{32}$ \\ergsec~\\citep{migliari2010}. The jet was found to be present in the source's hard state~\\citep{vanstraaten2000} comparable to what is now found in stellar mass black bole accretors and active galactic nuclei.  This is even more interesting as many indications place \\b0614 into a rare class of ultra-compact binaries with orbital periods P $\\approxlt 80$ min considered to be near the minimum period for LMXBs with hydrogen-rich main-sequence donors. \\citet{shahbaz2008} favor a 51.3 min binary period which would further classify this system as a neutron star accreting either from a hydrogen-deficient star or a degenerate, likely white dwarf companion~\\citep{verbunt1995}. Observations with \\chandra determined Ne K shell absorption significantly above the expectation from ISM column densities~\\citep{paerels2001}. This seems to support the suggestion by~\\citet{juett2001} and ~\\citet{juett2005} that the enhanced Ne/O ratios determined from X-ray spectra point to a low-mass, neon-rich degenerate white dwarf companion in \\b0614 and other very low period binaries and that these binaries are all ultra-compact. VLT spectroscopy confirms the deficiency of H and He in the accretion disk of the proto-type ultra-compact source \\u1626 and in \\b0614, again pointing towards an ultra-compact nature~\\citep{werner2006}. However, these data show a peculiar lack of Ne line which seems in stark contrast to the X-ray findings and it was suggested that the donor could be an eroded C/O white dwarf with no excessive Ne overabundance.  We observed \\b0614 with the high resolution transmission grating spectrometer (HETGS) onboard \\chandra for about 2.5 days in order to study in depth the X-ray absorption properties. In this paper we specifically focus on the intrinsic source properties, i.e. the nature of its X-ray continuum emission, possible broad line emissions, excess K edge properties, and spectral variability. ",
        "conclusions": "Our analysis shows that \\b0614 is observed in a hard state. The continuum is dominated by a powerlaw of index of about 2.2 similar to the one suggested by ~\\citet{piraino1999} and very recently by ~\\citet{madej2010}. We also confirm the need for a hard blackbody of about 1.2 keV as suggested by ~\\citet{piraino1999} and ~\\citet{migliari2010}. However, the spectra do not show a dominant disk blackbody component as found by ~\\citet{migliari2010} in \\swift data. On the other hand, we searched for an Fe K edge as suggested by ~\\citet{piraino1999} indicating the existence of a highly absorbed high energy component to no avail.  \\b0614 has been classified as an ultra-compact binary system. Specifically the detection of an excessive Ne K edge in an early \\chandra observation~\\citep{paerels2001} sparked interpretations that in combination with a possible ultra-short orbital period the donor stars are C-O or O-Ne-Mg stars~\\citep{juett2001}. Our observations indeed confirm the existence excess Ne K edge optical depth. A recent VLA study on the other hand finds a rather different account and observed Ne~II and~III emissions which suggested Ne to be underabundant ~\\citep{werner2006}. Our new observations show that the Ne K edge clearly shows an excess source intrinsic component. The edge is highly variable which is evidence that the excess Ne is source intrinsic. We conclude on this basis that \\b0614 indeed shows a Ne excess abundance. Whether this overabundance is restricted to Ne alone is uncertain. In at least one observation we find excess Mg as well. For other elements we do not find any conclusive overabundance signatures.  Another aspect of our study is the dynamic appearence of the edge. We find a smeared edge which we approximate by blue- and redhsifted components. To our knowledge this is the first time this kind of dynamics has been observed in a cool K edge. We tested this dynamics with known periodicities~\\citep{shahbaz2008} and conclude that it is random in nature. The Ne edge shows systematic blue- and red-shifts at average Doppler velocities of around 3500 \\kms. If we associate this cool absorbing material with an extended outer edge rim of the accretion disk and near edge-on view we can infer an upper limit to the outer radius of the disk of less than 10$^9$ cm (or $<$3500 \\kms) consistent with an ultra-compact binary. In order to make this feasibkle we have to assume that the disk radius is small enough in order to have X-ray with blue- and red-shifted absorption scattered in our line of sight. This may not be unfeasible but needs future detailed modeling. The random appearance of optical depth variations as well as edge shifts may also signal the presence of heavy turbulence of cool plasma entering the accretion resulting in heavy convulsions producing red- and and blue-shifts at times.  There is at least one broad line emission line complex in the spectra near ionized oxygen, which has already been seen in several previous observations with \\exo~\\citep{singh1994}, \\ros~\\citep{schulz1999}, \\sax~\\citep{piraino1999}, and \\swift~\\citep{migliari2010}. In neither of these obervations could the line be resolved because of insufficient instrument resolution. Even though the HETG in our study fully resolves the line, it still finds it broad on previously suggested levels and does not show structure. We find the width between 50 and 66 eV similar to previous studies. Very recent \\xmm observations also find a single broad line in RGS spectra \\citep{madej2010} and it is quite obvious that the line feature cannot be further resolved.  The line location suggests that it is dominated by \\oviiia\\ with maybe some contribution of \\oviib~ and \\oviiib~ emission. A single gaussian line interpretation leads to Doppler velocities of the order of 3$\\times10^4$ \\kms, which is 10$\\%$ of the speed of light and here we need to include relativistic effects such as gravitational (GR) broadening. \\citet{madej2010} successfully fit a GR profile to the line providing a line centroid of 18.5\\AA, an inner radius of 3.5 R$_g$, and a line of sight inclination of 88$^{\\circ}$. Our fits to the four observation segments give similar results, however also produce a larger spread of values. The inclination angles appears to be similar, an averaged value gives $87.5^{\\circ}$.  If these broad lines originate from the accretion disk, a viewing angle near 90$^{\\circ}$ makes sense and is already argued for in ~\\citet{migliari2010}. In general, disk velocities of 3$\\times10^4$ \\kms~ already require disk radii of a few $10^7$ cm and significantly lower inclination values would push the emission site beyond the inner edge of the disk. With a source luminosity of 0.02 L$_{edd}$ this requires ionization parameters far beyond $10^4$ and consequently barely allow for ionization fractions for H-like ions. We have observed high ionization parameters in the high flux state of Cir X-1, where edge-on viewing also seems likely~\\citep{brandt2000, schulz2008}. However, there the higher luminosities produce measurable ionization fractions at much larger disk radii. In the case of \\b0614 lower inclinations thus would lead to unrealistic dynamics, emission radii, and ionization parameters. We see some resemblance to the case of \\u1626~\\citep{schulz2001, krauss2007}, which being viewed face on already revealed disk (v sin i) velocities of several thousand \\kms. \\u1626 viewed edge-on would also likely show high velosity broadening in its lines. Recently the presence of soft broad line emission has also been reported in the Z-sources GX 349+2~\\citep{iaria2009} and Cyg X-2~\\citep{schulz2009}.  An identification of \\oviiia would also imply blue-shifts larger than 0.5~\\AA\\ independent of the line model. Even though the gravitational model predicts some blue-shift~\\citep{laor1991}, it should be consistent in all observations. However, the observed line wavelengths show significant variations indicating some additional dynamics in the line emitting region. Furthermore, \\chandra lines show consistently higher blus-shifts compared to the \\xmm measurement~\\citep{madej2010}. Possibilities may include changes in ionization balance, variable Compton scattering, or inner disk outflows. Changes in ionization balance are not obvious as observed spectral and flux changes do not appear dramatic. Compton upscattering requires very high optical depths as can be found in the disk but should lead to some tail on the blue side of the line which is not obvious in our data. An outflow interpretation remains viable, but needs more detailed modeling in order to quantify further. The difference in shift is about 0.7~\\AA\\ corresponding to an order of $10^4$ \\kms. Observed outflow velocities of this amount are rare in Galactic binaries, though. Winds and outflows in X-ray binaries are now commonly observed, mostly in systems exhibiting strong X-ray luminosities, such as in \\cirx1~\\citep{brandt2000, schulz2002a}, GX 13+1~\\citep{ueda2004} or the blackhole binary GRS J1655-40~\\citep{miller2006}, but all at much lower velocities. However, strong outflows of several 1000 \\kms\\ now also have been observed in sub-Eddington accretors such as 4U~1822-37~\\citep{bayless2010} and Her X-1~\\citep{boroson2007} in UV spectra. \\b0614 has gained a reputation in recent years of being of specifically violent nature showing sporadic abnormal type I superbursts~\\citep{kuulkers2009} and an active jet during its hard state quite similar to what is commonly observed in black hole binaries~\\citep{migliari2010}.  There are also several issues with respect to the observed abundances in the spectra. If \\b0614 is indeed ultra-compact we would expect that the donor star is either a C-O or O-Ne or similar white dwarf. In this respect it makes sense that we observe prodominantly neon and oxygen-rich matter in the accretion disk. However, some of the findings raise more questions. First of all, if we indeed see highly ionized oxygen in the inner disk, why are there no similar \\neix\\ and \\nex\\ emissions. The emissivity distributions with respect to temperature of these ions strongly overlap and, like in the case of \\u1626, we whould observe similar strong emissions from O and Ne. Our spectra allow for broad line flux upper limits at \\neix\\ and \\nex\\ wavelengths which are over an order of magnitude fainter than what we observe in \\oviiia\\ which does not provide an answer to the problem. Likewise, our analysis of the K edges reveals significant excesses in the Ne K columns, but not in the O K columns. At radii of the order of $10^9$ cm and the observed source luminosities we do not expect ionization parameters high enough to destroy all neutral oxygen.  Even though we lack answers to many of our findings, we tentatively conclude the following: \\begin{itemize} \\item{\\b0614 shows a variable excess Ne absorption column densities of up to several 10$^{18}$ cm$^{-2}$} \\item{The overabundance has a dynamic high velocity component and is evidently source intrinsic} \\item{The Ne K edge velocity smear could point to a probable outer disk radius limit of several $10^9$ cm supporting an ultra-compact binary nature} \\item{The soft X-ray line emission appears extremely broad with corresponding velocities of 3$\\times10^4$ \\kms and likely originates from a relativistic disk regime.} \\item{The line emission indicates an ionized disk layer, predominantly \\oviiia, at radii well below $10^8$ cm down to only several gravitational radii.} \\item{The soft X-ray line emission shows variable blue-shifts with respect to the \\oviiia rest wavelength indicating further dynamical processes near the inner disk.} \\end{itemize}"
    },
    "1005/1005.3631_arXiv.txt": {
        "abstract": "The use of standard rulers, such as the scale of the Baryonic Acoustic oscillations (BAO), has become one of the more powerful techniques employed in cosmology to probe the entity driving the accelerating expansion of the Universe. In this paper, the topology of large scale structure (LSS) is used as one such standard ruler to study this mysterious `dark energy'. By following the redshift evolution of the clustering of luminous red galaxies (LRGs) as measured by their 3D topology (counting structures in the cosmic web), we can chart the expansion rate and extract information about the equation of state of dark energy. Using the technique first introduced in (\\cite{PK}), we evaluate the constraints that can be achieved using 3D topology measurements from next-generation LSS surveys such as the Baryonic Oscillation Spectroscopic Survey (BOSS). In conjunction with the information that will be available from the Planck satellite, we find a single topology measurement on 3 different scales is capable of constraining a single dark energy parameter to within $5\\%$ and $10\\%$ when dynamics are permitted. This offers an alternative use of the data available from redshift surveys and serves as a cross-check for BAO studies. ",
        "introduction": "In the late 90's, the expansion of the Universe, first detected by Hubble, was confirmed from the precise measurements of SNIa and astonishingly found to be accelerating (\\cite{Riess1998,Perlmutter:1998np,Lange:2000iq,Hoekstra:2002xs,Riess:2004nr,Cole:2005sx,Astier:2005qq,WMAP3,Riess:2006fw}). This provided convincing evidence for the presence of an unidentified entity in the Universe which (in the context of General Relativity) must act against the gravitational attraction of ordinary matter.  Indirect yet also compelling evidence came from the missing energy density inferred from the discrepancy between the measurements of matter density (from direct measurements and observation of the Integrated Sachs-Wolfe effect) and the indications of spatial flatness from the CMB anisotropy spectrum, as well as the level of the initial inhomogeneity measured in the CMB compared with large scale structure today. Although the presence of this `dark energy' is now well-established, its nature still evades us and characterizing it has become one of the most important topics in cosmology today.   This is evidenced by the large number of experiments which have been proposed and designed with this question in mind.   One such effort is the Baryon Oscillations Spectroscopic Survey (BOSS) which plans to map the spatial distribution of luminous red galaxies (LRG) and quasars over $10,000$ sq. deg. of the sky. The survey hopes to detect the excess of galaxy clustering at $100$ Mpc/h separations left over from the acoustic oscillations in the baryon distribution at the time of last scattering.  The change in the characteristic scale of this phenonemon from the time of the CMB to today is encapsulated by the diameter angular distance $d_A(z)=(1+z)^{-1} r(z)$, which is related to the expansion rate of space via the comoving distance \\be r(z) = c \\int^{z}_0 \\frac{dz'}{H(z')}. \\label{rz} \\ee where \\be H(z) = H_0 \\sqrt{\\Omega_m\\bra{1+z}^3+\\Omega_{X} \\exp \\bra{\\int_0^z \\frac{1+w(z')}{1+z'} dz'}} \\label{w} \\ee where $H_0$ is the Hubble parameter today (note that flatness is assumed), $c$ is the speed of light, $\\Omega_m$ is the current matter density and $\\Omega_X$ is the current density of dark energy. We focussing on measuring the equation of state $w(z)$ of the dark energy component, which describes the ratio of its pressure to its energy density as a function of redshift. Because the scale of the oscillations at the time of last scattering is measured precisely from the CMB peak morphology, the BAO scale becomes a standard ruler. BOSS is forecasted to measure $d_A(z)$ to $1\\%$ at various redshifts (\\cite{BOSS}), placing constraints on the equation of state of dark energy $w$ in Eqn.~ \\ref{w}. In this paper, we will use the topology of large scale structure as another such standard ruler with which to get a handle on $w(z)$.  Being a tracer of the primordial density perturbations, LSS provides a record of the initial conditions and so its topology has been used extensively to test our current assumptions of the earliest epoch. For example, the measured topology of the Sloan Sky Digitial Survey (SDSS) traced by LRGs was shown to be consistent with expectations from a Universe with Gaussian randon phase initial conditions (\\cite{Gott_SDSS}).  In this paper, we recognize that topology is another measure of clustering or the number of structures per unit volume and thus by mapping how galaxy clustering per unit volume evolves with redshift, the expansion history of the Universe can be studied. The potential of this dataset in dark energy studies was first recognized in (\\cite{PK}), in which the genus statistic used to characterize LSS is adapted to map $r(z)$ at various epochs, thereby probing $w(z)$. Using the method presented in (\\cite{PK}), we build on this analysis with the aim of determining how effective the measurement of topology from future LSS surveys such as BOSS will be at elucidating the nature of dark energy.  We find that the topology is indeed effective in placing constraints of the average equation of state (at the $5\\%$ level). The constraints do however weaking when a smoothly evolving  equation of state is considered, with very little information on a second dark energy parameter delivered. In Section \\ref{genus_stats}, the genus statistics are reviewed, followed by a discussion of their application in Cosmology in Section \\ref{cosmo}. Section \\ref{analysis} gives the details of the analysis. In Section \\ref{results}, the method is applied to the BOSS data as well as information about the matter density and Hubble parameter from the Planck satellite. ",
        "conclusions": "\\label{conclusion} The question of the nature of dark energy is one of profound importance in cosmology today. The currently favoured model proposes that it is energy density associated with the vacuum, with a constant EOS of $w=-1$, identified as the mathematical equivalent of cosmological constant $\\Lambda$ in the Einstein field equation. In the context of our current theories of structure formation, $\\Lambda$ appears to be most successful in reproducing a wide range of present-day observations. However, it faces serious theoretical opposition \\emph{and} several recent works have highlighted the biases towards $w=-1$ that could potentially arise from the inclusion of \\emph{a prior} information when fitting the current data (\\cite{bassett-2004-617, Linder:2004rm}). Future experiments offer the exciting prospect of determining whether the cosmological constant is indeed the correct model.  BOSS is forecasted to place constraints on $w$ at the level of a few percent using clustering of LRGs on the scale of $100$ Mpc.  \\cite{PK} proposed a way to use the nature of LRG clustering on a wider range of scales as measured by the 3D topology to probe the expansion rate.   Because genus statistics relate to the overall shape of the power spectrum as opposed to a single scale, they are robust against fortuitous noise measurements. Furthermore, the toplogy measurements offer an alternative use of the data from redshift surveys and may be used as a cross check of the conclusions drawn from other techniques, such as the BAO method which uses the oscillations in the matter power spectrum to extract information.  In this paper, we have evaluated the constraints on $w(z)$ that are achievable from the BOSS survey using the method presented in (\\cite{PK}). We found that the BOSS survey using the topology method alone is capable of placing constraints of $5\\%$ on a constant equation of state. Allowing for the possibility of time variation in $w(z)$ degrades the constraints on $w_0$ to $12\\%$, while providing weak evidence for a second dark energy parameter. Although the topology measurement of BOSS may not be capable of testing for dynamical behaviour in $w(z)$, it has been shown to provide a robust measurement of the average equation of state, which may be sufficient to rule out non-$\\Lambda$ models."
    },
    "1005/1005.4227_arXiv.txt": {
        "abstract": "We report the results of multicolour observations of 30 E/S0 galaxies with dust lanes. For each galaxy we obtained broad-band images and narrow-band images using interference filters isolating the H$\\alpha$+[NII] emission lines to derive the amount and morphology of dust and ionized gas. To improve the wavelength coverage we retrieved data from the SDSS and 2MASS and combined these with our data. Ionized gas is detected in 25 galaxies and shows in most cases a smooth morphology, although knots and filamentary structure are also observed in some objects. The extended gas distribution closely follows the dust structure, with a clear correlation between the mass of both components. An extinction law by the extragalactic dust in the dark lanes is derived and is used to estimate the dust content of the galaxies. The derived extinction law is used to correct the measured colours for intrinsic dust extinction and the data are fitted with a stellar population synthesis model. We find that the H$\\alpha$ emission and colours of most objects are consistent with the presence of an ``old'' stellar population ($\\sim10$ Gyr) and a small fraction of a ``young'' population ($\\sim10-100$ Myr). To check this we closely examine NGC5363, for which archival Spitzer/IRAC and GALEX data are available, as a representative dust-lane E/S0 galaxy of the sample. ",
        "introduction": "For many years elliptical galaxies (Es) were considered as gas-poor objects (see Schweizer 1987 for a review), which formed in a single burst and evolved with no further star formation (Eggen et al.\\ 1962). However, improved detector sensitivity along with new discoveries of gas and dust of low column density in E/S0s brought fresh interest to this field. We now know that E/S0s host a multi-phase gas component with hot, X-ray emitting gaseous halos (Forman, Jones \\& Tucker 1985; Canizares Fabbiano \\& Trinchieri 1987; O'Sullivan, Forbes \\& Ponman 2001), warm ionized gas (IG) (Demoulin-Ulrich, Butcher \\& Boksenberg 1984; Phillips et al.\\ 1986), cold neutral-hydrogen (e.g., Knapp, Turner \\& Cunniffe 1985) and molecular gas  (e.g., Lees et al.\\ 1991).  The gas component in E/S0s is often morphologically associated with the presence of dust. Evidence for the presence of dust includes optical observations of obscured regions (Van Dokkum \\& Franx 1995; Ferrari et al.\\ 1999; Tomita et al.\\ 2000; Tran et al.\\ 2001) and the detection of far-IR emission from cold dust in a large fraction of E/S0s (Knapp et al.\\ 1989; Roberts et al.\\ 1991). The galaxies where dust obscuration produces well-defined dark lanes are known in the literature as dust-lane E/S0s.  The growing number of E/S0s in which a multiphase ISM is detected implies that the stellar population and abundances in this class of galaxies trace complex and varied star formation histories (SFHs). A recent burst of star formation in an E/S0 galaxy is expected to have a significant influence on the broad-band colours of the galaxy, even if the mass fraction of the young population is relatively small (see Fig.\\ 4 in Schawinski et al.\\ 2007). A recent analysis of optical spectra and of UV photometry in E/S0s shows that a significant fraction of these galaxies experienced low-level star formation within the past few Gyr  (e.g., Yi et al.\\ 2005; Di Matteo et al.\\ 2007; Kaviraj et al.\\ 2007).  The star formation activity in a certain galaxy can also be examined by studying emission-lines. The systematic search of IG in E/S0s with long-slit spectroscopic surveys allows the detection of various emission lines linked with excited and ionized gas (Caldwell 1984; Phillips et al.\\ 1986; Rampazzo et al.\\ 2005; Sarzi et al.\\ 2006; Serra et al.\\ 2008). Emission-line gas imaging enables the detecting of low levels of star formation even in faint galaxies with relatively modest integration times on small telescopes, and is therefore often used to reveal the full spatial extent of the IG. In particular, observing a galaxy with a filter centered at the rest-frame H$\\alpha$ line can, under reasonable assumptions, provide information on the amount of ionizing radiation emitted by stars. Surveys focusing on H$\\alpha$ emission imaging (Kim 1989; Shields 1991; Trinchieri \\& di Serego Alighieri 1991; Buson et al.\\ 1993; Goudfrooij et al.\\ 1994a; Singh et al.\\ 1995; Macchetto et al.\\ 1996; Martel et al.\\ 2004) detect IG in $\\gtrsim50$\\% of the E/S0s, and reveal various IG morphologies such as flattened disk-like or ring-like structures, or a filamentary structure.  This paper studies a sample of E/S0s where dust lanes (DLs) have been reported in the literature and focuses on the dust and IG components and their interrelation, and on the stellar population within the gaseous disks. The paper is organized as follows: \\S~\\ref{S:Obs_and_Red} gives a description of all the observations and data reduction, \\S~\\ref{S:results} details the broad-band and narrow-band flux calibration, and the data analysis is presented in \\S~\\ref{S:analysis}. The results are discussed in \\S~\\ref{S:discuss} and a summary of our conclusions is given in \\S~\\ref{S:conclude}. ",
        "conclusions": "\\label{S:conclude} We reported the properties of dust and ionized gas in a sample of 30 dust-lane E/S0 galaxies based on multicolour and H$\\alpha$ narrow-band imaging observations. For most galaxies, extended physically-associated dust and ionized gas structures are observed. The correlation between the dust and ionized gas is further illustrated by the correlation between the masses of these two components.  The extragalactic extinction law by the dust in the sample galaxies indicates similar properties to those canonical Galactic grains. This property is used to correct the flux measured within the ionized gas aperture for galactic extinction. The derived colours and H$\\alpha$ fluxes are fitted with stellar population synthesis models, with results indicating the presence of a small fraction of $\\sim10-100$ Myr old stars within the old stellar bulge.  We also discussed in detail the case of NGC5363 where we detected a unique spiral-like gas and dust distribution. The ISM characteristics combined with the wide spectral coverage available for this galaxy further support the late star-forming episode often observed in E/S0s. \\subsection*"
    },
    "1005/1005.5198_arXiv.txt": {
        "abstract": "To assess how future progress in gravitational microlensing computation at high optical depth will rely on both hardware and software solutions, we compare a direct inverse ray-shooting code implemented on a graphics processing unit (GPU) with both a widely-used hierarchical tree code on a single-core CPU, and a recent implementation of a parallel tree code suitable for a CPU-based cluster supercomputer.  We examine the accuracy of the tree codes through comparison with a direct code over a much wider range of parameter space than has been feasible before. We demonstrate that all three codes present comparable accuracy, and choice of approach depends on  considerations relating to the scale and nature of the microlensing problem under investigation. On current hardware, there is little difference in the processing speed of the single-core CPU tree code and the GPU direct code, however the recent plateau in single-core CPU speeds means the existing tree code is no longer able to take advantage of Moore's law-like increases in processing speed. Instead, we anticipate a rapid increase in GPU capabilities in the next few years, which is advantageous to the direct code. We suggest that progress in other areas of astrophysical computation may benefit from a transition to GPUs through the use of ``brute force'' algorithms, rather than attempting to port the current best solution directly to a GPU language -- for certain classes of problems, the simple implementation on GPUs may already be no worse than an optimised single-core CPU version. ",
        "introduction": "\\label{sct:intro}  Gravitational microlensing is the deflection of light by matter in the regime where multiple-imaging and high magnification events occur, but the angular separation of images is too small to resolve.  The observable effect is a change in brightness of a source over a period of time, with the details of the light curve depending on the complexity of the lens distribution and the surface brightness profile of the source. Gravitational microlensing has been used to study high magnification events on local scales, viz. the searches for dark matter compact objects in the Galactic bulge and halo (see for example references in and citations to \\citealt{alcock+93}; \\citealt{aubourg+93}; \\citealt{udalski+93}; \\citealt{calchi+02}; \\citealt{wyrzykowski+09}) and extrasolar planets (e.g. \\citealt{gould+92}; \\citealt{bond+04}; \\citealt{sumi+10}), and quasar microlensing on cosmological scales (e.g. \\citealt{vanderriest+89}; \\citealt{irwin+89}; \\citealt{k04}; \\citealt{bate+08}; \\citealt{dai+10}). For recent reviews of microlensing on all scales, see \\citet{wambsganss06}, \\citet{kochanek+07}, \\citet{gould08}, \\citet{mao08} and \\citet{gaudi10}.  Microlensing on local scales occurs in the low optical depth regime, where the source is affected by at most a few microlenses along the line of sight. Large numbers of stars must therefore be monitored in order to observe a single microlensing event. Conversely, the microlensing optical depth at the location of lensed quasar images is likely to be $\\sim1$. This microlensing signal is therefore more complex, as it is the result of a large number of microlenses acting on the source essentially all the time \\citep{wambsganss06}. In the remainder of this work, we focus on high optical depth microlensing.  A key element in many microlensing calculations is the determination of the magnification map; a typical example is shown in Figure \\ref{fig:map}. This map provides an estimate of the point-source brightening over a finite region of a two-dimensional source plane.  The magnification of an extended source is found by convolving a source profile (intensity and geometry) with the magnification map.  This approach has been used to probe the structure of quasar emission regions, including accretion discs (e.g. \\citealt{k04}; \\citealt{eigenbrod+08}; \\citealt{floyd+09}) and broad emission line regions (e.g. \\citealt{lewis+04}; \\citealt{abajas+07}).  \\begin{figure} \\centering \\scalebox{0.48}{\\includegraphics{fig1}} \\caption{A sample microlensing magnification map for a microlensing model with the following parameters: total convergence $\\sigma = 0.4$, shear $\\gamma=0.0$, $N_{\\rm pix} = 1024^2$ pixels in the source plane and $N_{\\rm rays} = 2000$ light rays per source pixel (on average). These parameters correspond to $N_{*} = 558$ lenses. The source plane has a side length of 20 Einstein radii. The greyscale covers the range from low magnification (black) to high magnification (white). } \\label{fig:map} \\end{figure}  The limitation in many microlensing calculations is the computational overhead in producing magnification maps with sufficient resolution and high (statistical) accuracy. Fitting observed light-curves obtained over long periods of time requires large maps with high pixel resolution to resolve small sources (such as quasar X-ray emission regions). Similarly, fitting multi-wavelength observations requires both large maps and high pixel resolution in order to simultaneously model emission from small optical emission regions and larger infrared emission regions. The situation is even worse when attempts are made to simultaneously model microlensing of accretion discs and the much larger broad emission line regions, a difference in scale of 2--3 orders of magnitude. Finally, many maps must be generated in order to ensure statistical independence of simulated measurements. It is only reasonable to sample any one magnification map as many times as you can place a source on it without overlapping.  In this paper we compare three techniques for computing magnification maps: a hierarchical tree code for single-core CPUs; a parallel, large data tree code for use on a CPU-based cluster supercomputer; and a direct, inverse ray-shooting code for use on graphics processing units (GPUs). We consider the accuracy of the codes, to determine whether the approximations inherent in the tree code are valid over a wider range of parameter space than has been feasible to test before, and compare run-times between the single-core tree code and the GPU code over a range of astrophysically-motivated parameters.  While the ``brute force'' approach to ray shooting has not been considered a feasible option on CPU architectures, we show that the GPU implementation is highly competitive, and is indeed faster than the tree code in many circumstances.   The expected increase in GPU speeds in the years ahead bodes well for advances in gravitational microlensing simulations. Our work hints at the potential benefit to other fields of astrophysical computation from the implementation of simple, ``brute force'' algorithms as a means of accelerating the adoption of GPUs in astronomy for time-consuming computations.  The remainder of this paper is set out as follows.  In Section \\ref{sct:magmap}, we describe the process of generating magnification maps, and summarise key features of three gravitational microlensing codes which can be used for this task.  In Section \\ref{sct:results}, we present a series of code comparisons (accuracy and speed). In Section \\ref{sct:discussion}, we discuss the relative merits of hierarchical and direct ray-shooting approaches, and suggest the classes of problems to which each is best-suited. We present our conclusions in Section \\ref{sct:conclusion}. ",
        "conclusions": "\\label{sct:conclusion} In this work, we have examined the accuracy of three highly competive alternatives for generating microlensing magnification maps: a widely-used hierachical tree code for single-core CPUs; a parallel, large data implementation for a distributed CPU-based computing cluster; and a brute force solution that utilises the massive parallelism of modern graphics processing units. We have demonstrated that these three codes do indeed present comparable accuracy, and choice of approach will depend on considerations relating to the scale and nature of the problem to be investigated.  While considered impractical as a calculating approach even a few years ago, our timing tests show that direct, inverse ray-shooting is now a viable, and indeed more accurate, option.  Moving a simple algorithm to GPU, rather than attempting to directly implement a more advanced solution -- originally developed to overcome a hardware restriction of single-core CPUs -- to achieve run-times that are already similar, suggests an approach that other early adopters of GPUs for astronomy computation may wish to consider."
    },
    "1005/1005.2895_arXiv.txt": {
        "abstract": "{} { We study an undocumented large translucent cloud, detected by means of its enhanced radiation on the SHASSA (Southern H-Alpha Sky Survey Atlas) survey. We consider whether its excess surface brightness can be explained by light scattered off the dust grains in the cloud, or whether emission from in situ ionized gas is required. In addition, we aim to determine the temperature of dust, the mass of the cloud, and its possible star formation activity. } { We compare the observed H$\\alpha$ surface brightness of the cloud with predictions of a radiative transfer model. We use the WHAM (Wisconsin H-Alpha Mapper) survey as a source for the Galactic H$\\alpha$ interstellar radiation field illuminating the cloud.  Visual extinction through the cloud is derived using 2MASS $J$, $H$, and $K$ band photometry. We use far-IR ISOSS (ISO Serendipitous Survey), IRAS, and DIRBE data to study the thermal emission of dust. The LAB (The Leiden/Argentine/Bonn Galactic HI Survey) is used to study 21\\,cm HI emission associated with the cloud. } { Radiative transfer calculations of the Galactic diffuse H$\\alpha$ radiation indicate that the surface brightness of the cloud can be explained solely by radiation scattered off dust particles in the cloud.  The maximum visual extinction through the cloud is about 1.2\\,mag. The cloud is found to be associated with 21\\,cm HI emission at a velocity $\\sim-9$\\,km\\,s$^{-1}$.  The total mass of the cloud is about 550--1000\\,M$_{\\sun}$. There is no sign of star formation in this cloud. The distance of the cloud is estimated from the Hipparcos data to be $\\sim 100$\\,pc. } {} ",
        "introduction": "del Burgo \\& Cambr\\'esy (\\cite{burgo06}) were the first to report on detection of diffuse H$\\alpha$ emission in a molecular cloud.  They detected an excess surface brightness over the cloud \\object{LDN1780} of intensity $\\sim$1--4 rayleigh (one rayleigh (R) being equivalent to $2.4\\,10^{-7}$\\,ergs\\,cm$^{-2}$\\,s$^{-1}$\\,sr$^{-1}$ at the wavelength of H$\\alpha$ emission).  They interpreted the surface brightness as a result of enhanced in situ cosmic ray ionization.  However, Mattila et~al.\\ (\\cite{mattila07}) showed that the H$\\alpha$ surface brightness observed in LDN1780 can be explained solely in terms of scattered H$\\alpha$ radiation.  In addition, Mattila et~al.\\ found several other molecular clouds that had excess H$\\alpha$ emission relative to the surroundings of the cloud. In some cases, a cloud was not detected in H$\\alpha$ images, although other physically similar clouds do.  These observations can be naturally explained in the framework of scattered H$\\alpha$ radiation by varying the proportions of general diffuse in situ H$\\alpha$ emission either in front of or behind the dust cloud.  While comparing the all-sky H$\\alpha$ and 100\\,$\\mu$m IRAS maps, we noticed a large cloud visible in both maps. The Galactic coordinates of the cloud are $l \\approx 294\\degr$, $b \\approx -24\\degr$.  The maximum excess surface brightness of H$\\alpha$ is about 2.4\\,R.  With a size of about 1.4\\degr$\\times$4.9\\degr, the cloud is the largest dust cloud visible in the light of scattered H$\\alpha$ emission known to us. Hereafter the cloud is called G294-24. ",
        "conclusions": "We have studied the general properties of an undocumented large, nearby (distance $\\sim100$\\,pc) translucent cloud. The cloud is without star formation, and thus illuminated only by the general interstellar radiation field.  We have measured a drop in the temperature of the dust particles of about 1.2\\,K between the edge and the center of the cloud, the minimum dust temperature being $\\sim$17.5\\,K.  The total mass of the cloud is about 550--1000\\,M$_{\\sun}$. The cloud is detected in the 21\\,cm HI spin-flip line, and the atomic hydrogen mass of the cloud is about 40\\,M$_{\\sun}$.  The excess of diffuse H$\\alpha$ surface brightness of the cloud over the background sky can be naturally explained as the general interstellar radiation field being scattered off the dust grains."
    },
    "1005/1005.3966_arXiv.txt": {
        "abstract": "The conventional derivation of the gamma-ray burst afterglow jet break time uses only the blast wave fluid Lorentz factor and therefore leads to an achromatic break. We show that in general gamma-ray burst afterglow jet breaks are chromatic across the self-absorption break. Depending on circumstances, the radio jet break may be postponed significantly. Using high-accuracy adaptive mesh fluid simulations in one dimension, coupled to a detailed synchrotron radiation code, we demonstrate that this is true even for the standard fireball model and hard-edged jets. We confirm these effects with a simulation in two dimensions. The frequency dependence of the jet break is a result of the angle dependence of the emission, the changing optical depth in the self-absorbed regime and the shape of the synchrotron spectrum in general. In the optically thin case the conventional analysis systematically overestimates the jet break time, leading to inferred opening angles that are underestimated by a factor 1.32 and explosion energies that are underestimated by a factor 1.73, for explosions in a homogeneous environment. ",
        "introduction": "Gamma-ray burst (GRB) afterglows are produced by jetted outflows slowing down from ultrarelativistic velocities to non-relativistic velocities. Synchrotron emission results when the blast wave interacts with the circumburst medium and electrons get shock-accelerated and are accelerated in magnetic fields that are also usually presumed to be generated at the shock.  The jet nature of afterglows is not observed directly (there are no spatially resolved jet images). Theoretically, the energetics of the fireball model require that the outflow be concentrated towards the observer, since a spherical outflow would imply unphysically high explosion energy on the order of the Solar rest mass ($\\backsim 10^{54}$ erg) instead of a more reasonable $\\backsim 10^{52}$ erg. Observationally, afterglow light curves at various frequencies show a break followed by a steepening of the curve. This can be explained as a \\emph{jet break}, beyond which relativistic beaming effects have decreased in strength to an extent where it becomes possible to distinguish between a jet and spherical outflow. A lack of emission from beyond the jet opening angle is now seen. A second effect that is theorized to occur at roughly the same time is a diminishing of synchrotron emission due to the onset of significant lateral expansion of the jet.  The conventional theoretical argument for the appearance of a jet break compares the \\emph{half opening angle} $\\theta_\\mathrm{h}$ of the jet to the fluid Lorentz factor $\\gamma$ directly behind the shock front. Once the latter has decreased to $\\gamma \\backsim 1 / \\theta_\\mathrm{h}$, beaming has sufficiently decreased and the absence of emission from angles beyond $\\theta_\\mathrm{h}$ becomes noticable. Since this ignores the basic shape of the synchrotron spectrum as well as electron cooling and the increase in optical depth due to synchrotron self-absorption, this suggests an \\emph{achromatic} break in the light curve.  However, achromatic jet breaks are relatively rare in the literature. Observations in the \\emph{Swift} era often do not show a clear jet break in the X-ray when the optical light curve does show a break (e.g. GRB990510, GRB030329, GRB060206). Also, the jet opening angle inferred from radio observations is often much larger than that from (earlier) X-ray or optical observations, which is again at odds with the expected achromaticity of the break.  Different explanations for the lack of truly achromatic breaks have been put forth in the literature. Staying close to the data, first of all \\cite{Curran2008} show that current observations do not actually rule out an achromatic break between X-ray and optical and that both single and broken power law fits are often consistent with the X-ray data. On the theoretical side, the afterglow jet has been expanded into a structured jet (e.g. \\citealt{Rossi2002}) or a superposition of multiple hard-edged jets with different opening angle. By now superpositions of up to three jets have been used to fit the data of GRB030329, with the opening angle explaining the radio light curve significantly larger than the others \\citep{vanderHorst2005}. An overview of afterglow jet structure and dynamics can be found in \\cite{Granot2007}.  Oddly enough, the analytical argument leading to an achromatic break has never been challenged. In this paper we present results from high resolution numerical jet simulations, coupled to an advanced synchrotron radiation code, showing that the standard analytical argument systematically overestimates the jet break time.  We explain the numerical set-up of the simulations in this paper in section \\ref{set_up_section}. In section \\ref{optically_thin_section} we study jet breaks while ignoring both self-absorption and cooling and find that the basic calculation linking jet break time and opening angle can be off by days. We derive a relation between jet break time and opening angle in the optically thin case and compare limb-brightening of the afterglow image at different spectral regimes.  In section \\ref{jetbreaks_full_section} we calculate the full synchrotron spectrum and compare jet breaks at different spectral regimes. We find that the break at radio frequencies is postponed compared to the jet breaks observed at frequencies above the self-absorption break frequency. Depending on opening angle and observer frequency, this time difference can be on the order of several days (over a factor 2 in jet break time), even though we did not add any novel radiation physics or make any nonstandard assumptions. The new aspects of our calculation are merely the accuracy of the radiation code and the numerical resolution of the fluid simulation. Self-absorption is fully treated using linear radiative transfer equations and local electron cooling times are numerically calculated through an advection equation. The difference in jet break characteristics can be understood from the fact that different regions of the jet provide the dominant contribution for different observer frequencies. We show images of the emission coefficient throughout the blast wave to visualise the underlying physics.  Up to section \\ref{jetbreaks_full_section} we assume that collimated outflow can be represented by a conic section from a 1D spherically symmetric simulation. We test this assumption, which implies that lateral spreading has little effect on the observed jet break, by performing a 2D simulation in section \\ref{2D_simulation_section}. Higher dimensional simulations are not the focus of this paper and we will only briefly discuss the consequences of lateral spreading. We end with a summary and discussion of our results in section \\ref{jetbreaks_summary_section}. ",
        "conclusions": "\\label{jetbreaks_summary_section} We have presented simulation results for high-resolution adaptive mesh simulations of afterglow blast waves in one and two dimensions. We have used explosion parameters that nominally apply to GRB 030329, but these were derived by earlier authors using a purely analytical model (in \\cite{vanEerten2009c} we have shown that the difference between analytical model and simulation is significant. A paper improving fit models using simulation results is currently in preparation). The aim of these simulations is to study the achromaticity of the afterglow jet break. We have found that the jet break is chromatic between frequencies below both the self-absorption break $\\nu_{a}$ and the synchrotron break $\\nu_m$ and above these breaks.  A result that is of direct practical use is an improvement on the relationship between observed jet break time and opening angle. We show how the conventional analysis for jet breaks in optically thin jets, systematically underestimate the jet opening angle. For a homogeneous medium, this amounts to a factor 1.3, but we give an expression for general circumburst medium density structure.  The difference in jet break time between optically thick and optically thin frequencies is explored in some detail. We have calculated emission profiles for the jet at various observer frequencies at ten days observer time. From these calculations it is shown that the region of the blast wave dominating the observed flux moves around for different frequencies. At high frequencies, this region lies more to the edge of the jet and to early emission times, whereas at high optical depth and frequencies below the synchrotron break the contributing region moves to the front and center of the blast wave -and therefore remains within the jet cone long after the high frequency region falls outside the jet cone.  The slope of the jet break will be different for different observer frequencies. If limb brightening is strong, the jet break slope for hard-edged jets (without lateral expansion) will initially overshoot its asymptotic value and then slowly evolve to its asysmptotic limit, which lies $(3-k)/(4-k)$ below the corresponding slope for the spherical case at the same observer time.  We have confirmed our results in two dimensions. The jet break in the radio is still postponed, but less so than for hard-edged jets. At 28 days in observer time, the effectively contributing region remains within the initial opening angle, even though lateral spreading of the fluid has already progressed noticably. As a result of the lateral spreading, the optical depth through the fluid decreases.  The chromaticity of the jet breaks is a result of the detailed interplay between the synchrotron radiation mechanisms and the fluid dynamics. Models that treat synchrotron radiation in a simplified manner, like \\cite{Zhang2009}, which does not include self-absorption and does not locally calculate the cooling times, will not reveal a chromatic jet break. \\cite{Granot2007} does not discuss self-absorption in the context of the jet break.  Finally we note that we have in this entire paper assumed the observer to lie precisely on the jet axis. Although the observer angle will be small in practice compared to the jet opening angle (in order to be able to observe early afterglow and prompt emission in the first place), this is nevertheless likely to have a profound effect on the shape and timing of the jet break. We are expanding the radiation code to deal with off-axis observers and will report the results in our follow-up study."
    },
    "1005/1005.3820_arXiv.txt": {
        "abstract": "{The {\\it Herschel} Lensing Survey (HLS) will conduct deep PACS and SPIRE imaging of $\\sim$40 massive clusters of galaxies. The strong gravitational lensing power of these clusters will enable us to penetrate through the confusion noise, which sets the ultimate limit on our ability to probe the Universe with {\\it Herschel}.  Here, we present an overview of our survey and a summary of the major results from our Science Demonstration Phase (SDP) observations of the Bullet Cluster ($z=0.297$).  The SDP data are rich, allowing us to study not only the background high-redshift galaxies (e.g., strongly lensed and distorted galaxies at $z=$2.8 and 3.2) but also the properties of cluster-member galaxies.  Our preliminary analysis shows a great diversity of far-infrared/submillimeter spectral energy distributions (SEDs), indicating that we have much to learn with {\\it Herschel} about the properties of galaxy SEDs.  We have also detected the Sunyaev-Zel'dovich (SZ) effect increment with the SPIRE data.  The success of this SDP program demonstrates the great potential of the {\\it Herschel} Lensing Survey to produce exciting results in a variety of science areas.} ",
        "introduction": "With the successful launch and commissioning of the ESA's {\\it Herschel} Space Observatory \\citep{Pilbratt10}, we are again on the verge of making great new discoveries.  Following the breakthrough submillimeter/millimeter observations with SCUBA and MAMBO \\citep[c.f.,][]{Blain02}, deep MIPS 24 $\\mu$m observations carried out by the {\\it Spitzer} Space Telescope have enabled us to trace the evolution of infrared-luminous galaxies up to $z \\sim$ 3--4 \\citep[e.g.,][]{Pgperez05,LeFloch05}.  However, the validity of all these {\\it Spitzer}-based results rests on the assumption that the total infrared luminosities of high-redshift infrared galaxies can be estimated accurately by sampling their rest-frame mid-infrared emission (e.g., the MIPS 24$\\mu$m band samples the rest-frame 8 $\\mu$m emission at $z=2$).  Indeed, some {\\it Spitzer} results have already questioned this assumption, suggesting that the use of local galaxy spectral energy distribution (SED) templates may lead to overestimating the total infrared luminosities of high-redshift galaxies \\citep[e.g.,][]{Papovich07,Rigby08}.  {\\it Herschel} will allow us to measure the total infrared luminosities of a large number of high-redshift galaxies directly for the first time.  With {\\it Herschel}, confusion noise produced by a sea of blended faint galaxies sets the ultimate limit on how deeply we can probe the Universe.  Once the source confusion sets in, it is no longer possible to improve the detection limit by integrating longer.  This limitation is especially severe for SPIRE, which reaches the confusion limit quickly.  To penetrate through the confusion limit, gravitational lensing by massive galaxy clusters offers a very powerful and yet cheap solution \\citep[e.g.,][]{Blain97}.  Magnification factors of 2--4x are quite common in the cluster core regions, and when a background source is strongly lensed (i.e., multiply imaged), magnification factors can reach 10x--30x or more.  Note that a magnification factor of 10x corresponds to a factor of 100x saving in observing time when the sensitivity is background-limited.  Therefore, a fairly short-integration image of a cluster core region would often reveal sources that are well below the detection limit of an ultra-deep blank-field image.  This method was for example employed for the first SCUBA observations of the high-redshift Universe, which resulted in the identification of the substantial population of infrared-luminous galaxies at $z>1$ \\citep{Smail97}.  The use of gravitational lensing is especially powerful at infrared/submillimeter wavelengths. This is because cluster cores are dominated by early-type galaxies, which usually emit little at these wavelengths.  Therefore, infrared/submillimeter sources detected in cluster cores are often background galaxies.  In other words, when observed at these wavelengths, massive cluster cores virtually act as a transparent lens.  Lensing studies in the infrared/submillimeter also benefits from the steep galaxy counts at these wavelengths.   These types of lensing surveys, however, have one limitation: the small number of strongly lensed galaxies observed per cluster. Although a large number of massive clusters have been targeted by ground-based submillimeter/millimeter observations \\citep[e.g.,][]{Smail02,Chapman02,Knudsen08}, the number of strongly lensed (i.e., multiply imaged) galaxies discovered in these surveys remains small: for example, the $z=2.5$ galaxy in Abell 2218 \\citep{Kneib04}, the $z=2.9$ galaxy in MS0451.6-0305 \\citep{Borys04}, the $z=2.8$ galaxy in the Bullet Cluster \\citep{Wilson08,Gonzalez09,Rex09,Johansson10}, and most recently the exceptionally bright $z=2.32$ galaxy in MACSJ2135-010217 \\citep{Swinbank10}.  Based on these observations, we empirically estimate the rate of finding such strongly lensed infrared/submillimeter galaxies to be roughly 1 in 10 with the sensitivity of ground submillimeter observations (e.g., SCUBA, LABOCA).  Therefore, many tens of clusters need to be observed to study a significant number of strongly lensed galaxies. ",
        "conclusions": "The SDP observations of the Bullet Cluster clearly demonstrate the great potential of the HLS in a variety of science areas.  With the {\\it Herschel} observations nearly done, HLS is expected to make a rapid progress in the near future. One immediate interest is whether it can find strongly lensed sources that are bright enough to perform spectroscopy with {\\it Herschel}.  Ultimately, we will construct a definitive sample of $\\sim$50 lensing clusters with a variety of multi-wavelength data and accurate mass models.  Such a data set can be further exploited by future facilities such as {\\it ALMA}, {\\it JWST}, {\\it SPICA}, and ground 30-meter class telescopes."
    },
    "1005/1005.1014_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1} {\\it RTS2} originates from {\\it RTS}, {\\it the Remote Telescope System}. {\\it RTS} was written during years 1999 and 2000 as a team project for Computer Science courses\\cite{jilek00} by students of Mathematics and Physics faculty of Charles University in Prague.  {\\it RTS2} is primarily developed on the {\\it Ubuntu Linux} distribution. It is known to run on a wide variety of Linux distributions, the {\\it Solaris} operating system and partially on {\\it Microsoft Windows} and {\\it Mac OS X}.  Project source code can be obtained from the Web site {\\tt http://rts2.org}. The project Wiki pages on {\\tt http://rts2.org/wiki} list group experiences with various fully autonomous observatory projects.  The outline of this paper is as follows: The first section is this introduction. The next section describes goals and history of the project development. The third section focus is a description of how various tasks inside {\\it RTS2} are divided.  The fourth summarises communication protocol. The next section describes approaches used during development. The sixth section forms the core of the article -- it lists currently added features. The seventh section informs readers about currently running {\\it RTS2} systems. The eighth is a reminder about our experiences with restarting failed devices. The contribution concludes with a section about expected developments. ",
        "conclusions": "\\label{sec:5} This article presented an open-source system for robotic observatory control. It provides an overview of rationales for its development, its composition and its main features. It then focused on a list of recent improvements. The article concludes with lists of development items still left on the agenda.  \\vskip 1cm {\\it Acknowledgement} The author would like to acknowledge generous financial support provided by Spanish {\\it Programa de Ayudas FPI del Ministerio de Ciencia e Innovaci\u00f3n (Subprograma FPI-MICINN)} and European {\\it Fondo Social Europeo}. Work on {\\it RTS2} was supported, influenced and encouraged by numerous people, whose list would be too large for this article. Persons from this list which according to author deserve to be mention explicitly are Martin Jel\\'inek, Alberto Castro-Tirado, Antonio de Ugarte Postigo, Ronan Cunniffe, Michael Prouza, Ren\\'e Hudec, Victor Reglero and Beatriz S\\'anchez F\\'elix. The article was carefully reviewed by two anonymous referees, whose suggestions significantly improved it. Stephen Bailey made final gramatical improvements to the article."
    },
    "1005/1005.2221_arXiv.txt": {
        "abstract": "We present an analysis of elemental abundances of He, N, O, Ne, S, and Ar in Magellanic Cloud planetary nebulae (PNe), and focus initially on 14 PNe in the Small Magellanic Cloud (SMC). We derived the abundances from a combination of deep, high dispersion optical spectra, as well as mid-infrared (IR) spectra from the \\textit{Spitzer Space Telescope}. A detailed comparison with prior SMC PN studies shows that significant variations among authors of relative emission line flux determinations lead to systematic discrepancies in derived elemental abundances between studies that are $\\gtrsim0.15$~dex, in spite of similar analysis methods. We used ionic abundances derived from IR emission lines, including those from ionization stages not observable in the optical, to examine the accuracy of some commonly used recipes for ionization correction factors (ICFs). These ICFs, which were developed for ions observed in the optical and ultraviolet, relate ionic abundances to total elemental abundances. We find that most of these ICFs work very well even in the limit of substantially sub-Solar metallicities, except for PNe with very high ionization. Our abundance analysis shows enhancements of He and N that are predicted from prior dredge-up processes of the progenitors on the AGB, as well as the well known correlations among O, Ne, S, and Ar that are little affected by nucleosynthesis in this mass range. We identified MG~8 as an interesting limiting case of a PN central star with a $\\approx3.5$~M$_\\sun$ progenitor in which hot-bottom burning did not occur in its prior AGB evolution. We find no evidence for O depletion in the progenitor AGB stars via the O-N cycle, which is consistent with predictions for lower-mass stars. We also find low S/O ratios relative to SMC \\ion{H}{2} regions, with a deficit comparable to what has been found for Galactic PNe. Finally, the elemental abundances of one object, SMP-SMC~11, are more typical of SMC \\ion{H}{2} regions, which raises some doubt about its classification as a PN. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.2826_arXiv.txt": {
        "abstract": "This article deals with the problem of the motion of stars in galaxies. By using the Newton's theory combined with a gravitational time dilatation for the weak gravitational field, it is possible to give a solution without using the dark matter. ",
        "introduction": "Since the observations of Zwicky in \\cite{Zwicky33}, we wonder why the motion of stars in galaxies do not corroborate the gravitational effects. The main explanation is the \\emph{dark matter} which is hypothetical matter of the standard cosmological model $\\Lambda$-CDM that is undetectable by its emitted radiation, but whose presence can be inferred from gravitational effects on visible matter \\cite{Papantonopoulos07}. However, such a theory raises some problems \\cite{BlanchetBis09}. There exists an alternative theory to the dark matter which is the \\emph{Modified Newtonian dynamics} MOND of Milgrom who has developed a change in the Newton's law  \\cite{Milgrom83,Milgrom83bis,Milgrom83Ter}. A comparison between these two approaches is developed in \\cite{Blanchet09} by Blanchet who has also proposed a new kind of dark matter: the dipolar dark matter \\cite{BlanchetBis09}.\\\\ The main goal of this paper is to provide another solution for the problem of the motion of stars in galaxies by using a new gravitational time dilation for the weak gravitational field combined with the Newton's theory.\\\\ The paper is organized as follows. The problem of the motion of stars in galaxies is recalled in Section \\ref{SecDM} and the solution involving a new gravitational time dilation for the weak gravitational field is addressed in Section \\ref{SecTD}. Finally, a conclusion is given in Section \\ref{SecConc} ",
        "conclusions": "\\label{SecConc} Even if the $\\Lambda$-CDM model is the main theory, we cannot rule out other possibilities, as long as we have no direct evidence of the existence of dark matter. The gravitational time dilatation for the weak gravitational field is one of the possibility to explain the motion of stars in galaxies."
    },
    "1005/1005.2960_arXiv.txt": {
        "abstract": "We report the first observation of an anisotropy in the arrival direction of cosmic rays with energies in the multi TeV region in the Southern sky using data from the IceCube detector. Between June 2007 and March 2008, the partially-deployed IceCube detector was operated in a configuration with 1320 digital optical sensors distributed over 22 strings at depths between 1450 and 2450 meters inside the Antarctic ice. IceCube is a neutrino detector, but the data are dominated by a large background of cosmic ray muons. Therefore, the background data are suitable for high-statistics studies of cosmic rays in the Southern sky. The data include 4.3 billion muons produced by downgoing cosmic ray interactions in the atmosphere; these events were reconstructed with a median angular resolution of 3 degrees and a median energy of $\\sim20$~TeV. Their arrival direction distribution exhibits an anisotropy in right ascension with a first harmonic amplitude of  $(6.4\\pm0.2~$stat$. \\pm 0.8~$syst$.)\\times10^{-4}$. ",
        "introduction": "Long-term observations of cosmic ray muons by underground experiments have demonstrated the presence of an anisotropy in the cosmic ray intensity up to a few hundred GeV~\\citep{nag}.  Recent underground and surface array measurements of cosmic rays by the Tibet Array~\\citep{ta}, Super-Kamiokande~\\citep{sk} and Milagro~\\citep{milagro} indicate that the anisotropy persists into the TeV range.  All of the TeV measurements were performed in the Northern hemisphere; so far, no such measurement has been performed covering the entire Southern hemisphere at median energies in the multi TeV region. With the deployment of the IceCube Neutrino Observatory at the South Pole, we have for the first time measured the anisotropy at TeV energies in the Southern sky. IceCube is primarily a neutrino detector but it is sensitive to the muons produced in downward-going cosmic ray air showers. The observatory provides high-statistics measurements of cosmic rays with median energy of $20$~TeV.  When completed in 2011, IceCube will comprise 5160 optical modules buried 1450 and 2450 meters below the surface of the polar ice sheet.  The modules are physically connected to the surface by electronic umbilical lines, or ``strings,'' with 86 strings in total~\\citep{dommb}. In this paper, we use cosmic ray data recorded by the detector in its 22-strings configuration (IC22) between June 2007 and March 2008 to produce the cosmic ray skymap of the Southern sky in the TeV range. ",
        "conclusions": "Using a high-statistics sample of downgoing cosmic rays with a median energy of $20$~TeV and a median angular resolution of $3^\\circ$, we presented the first  map of the relative intensity of the flux of TeV cosmic rays in the Southern sky.  The arrival direction distribution of the cosmic rays is found to be anisotropic with a first harmonic amplitude and phase of $A_{1}=(6.4\\pm0.2~$stat$. \\pm 0.8~$syst$.)\\times10^{-4}$ and $\\phi_{1} = 66.4^\\circ\\pm2.6^\\circ~$stat$. \\pm 3.8^\\circ~$syst$.$. The observation appears to be a continuation of a previously-measured cosmic ray anisotropy reported in the Northern hemisphere~\\citep{ta,sk,milagro}.  The origin of the anisotropy remains unclear. Compton-Getting~\\citep{cg} suggested that the relative motion of the solar system around the Galactic center in the cosmic ray plasma should give rise to an excess in the direction of motion of the solar system and a deficit in the opposite direction. In this model, an excess flux should appear with a maximum in right ascension between $290^{\\circ}$ and $340^{\\circ}$ and a minimum in right ascension  between $110^{\\circ}$ and $160^{\\circ}$~\\citep{ta}. As shown in Fig 4, the excess can not be described in terms of the direction of motion of the solar system.  Therefore, we conclude that the Compton-Getting effect could be (at most) one of several contributions to the cosmic ray anisotropy. This effect will be addressed in more detail in a future study including the energy dependence of the anisotropy.   It is tempting to try to interpret the cosmic ray excess as an artifact of the heliospheric magnetic field. However, the maximum gyro-radius of a 10 TeV cosmic ray proton in a 1 $\\mu$G magnetic field is about 0.01 pc, i.e. much larger than the size of the heliosphere. As a consequence, the observed anisotropy is more likely to be connected to features of the local interstellar magnetic field at distances $<$ 1 pc. We are also investigating the possibility that the cosmic ray excess is associated with structures in the Galactic magnetic field at larger distance scales, or with diffusive particle flows from a nearby Galactic source such as Vela.  The still growing IceCube observatory will be completed in 2011 with a total of 86 strings and a  volume of 1 km$^3$. The estimated rate of cosmic ray-induced muons will be greater than  30 billion events per year. Such high statistical power, together with an estimated energy resolution of about 0.3 in log(E), will allow us to determine, in one year, the variation of cosmic ray anisotropy in several energy ranges up to a few hundred TeV. The energy dependence study will provide fundamental  hint at the nature of the source or sources of the cosmic rays, as well as their propagation through the Galactic magnetic field."
    },
    "1005/1005.1615_arXiv.txt": {
        "abstract": "{By means of different physical mechanisms, the expansion of H\\,{\\sc{ii}} regions can promote the formation of new stars of all masses. RCW~120 is a nearby Galactic \\HII\\ region where triggered star formation occurs. This region is well-studied - there being a wealth of existing data - and is nearby. However, it is surrounded by dense regions for which far infrared data is essential to obtain an unbiased view of the star formation process and in particular to establish whether very young protostars are present.} {We attempt to identify all Young Stellar Objects (YSOs), especially those previously undetected at shorter wavelengths, to derive their physical properties and obtain insight into the star formation history in this region.} {We use $\\it{Herschel}$-PACS and -SPIRE images to determine the distribution of YSOs observed in the field. We use a spectral energy distribution fitting tool to derive the YSOs physical properties. } {$\\it{Herschel}$-PACS and -SPIRE images confirm the existence of a young source and allow us to determine its nature as a high-mass (8-10~M$_{\\odot}$) Class~0 object (whose emission is dominated by a massive envelope M$_{env}$ $\\simeq$ 10$^{3}$~M$_{\\odot}$) towards the massive condensation~1 observed at (sub)-millimeter wavelengths. This source was not detected at 24\\,$\\mu$m and only barely seen in the MISPGAL 70\\,$\\mu$m data. Several other red sources are detected at $\\it{Herschel}$ wavelengths and coincide with the peaks of the millimeter condensations. SED fitting results for the brightest $\\it{Herschel}$ sources indicate that, apart from the massive Class~0 that forms in condensation 1, low mass (0.8\\,-- 4~M$_{\\odot}$) stars are forming around RCW~120 with ages younger than 5$\\times$10$^4$ years. This indicates that YSOs observed on the borders of RCW~120 are younger than its ionizing star, which has an age of about 2.5~Myr.} {$\\it{Herschel}$ images allow us to detect new YSOs that are too young and embedded to be detected at shorter wavelengths (25 of the 49 $\\it{Herschel}$ sources are new detections). This offers a new and more complete view of the star formation in this region. PACS and SPIRE fluxes were obtained for the brightest YSOs and allow us to strongly constrain both their spectral energy distribution and their physical properties through SED fitting. A more accurate determination of their properties allows us, for the first time, to discuss the star formation history in this region by comparing $\\it{similar}$ sources at $\\it{different}$ evolutionary stages. } ",
        "introduction": "RCW~120 (Rodgers et al.~\\cite{rcw60}) is a  bubble-shaped Galactic \\HII\\ region, located 1.3~kpc from the Sun. This region was studied in detail by Zavagno et al.~(\\cite{zav07}, hereafter ZA07) and Deharveng et al.~(\\cite{DE09}, hereafter DE09). ZA07 observed the cold dust emission at 1.2~mm in this region with SIMBA at SEST and found five condensations that border the southern half of the ionized region. Three other condensations are observed away from the region (see Fig.~\\ref{Fig1}). However, as is clearly observed in extinction in the $K_{\\rm{S}}$-band at 2.17\\,$\\mu$m (see Fig.~3 in ZA07), the region is surrounded by a shell of dense material. Dark filaments are observed in the northern part of the region in the $\\it{Spitzer}$-GLIMPSE 8\\,$\\mu$m image (see Fig.~3 in ZA07 and Fig.~\\ref{Fig1}). ZA07 also used the 2MASS and $\\it{Spitzer}$-GLIMPSE point source catalogs to study the nature and spatial distribution of the YSOs observed towards this region. They found numerous Class~I and Class~II sources towards the millimeter condensations that border the ionized region. The association of these young stellar objects (YSOs) with the RCW~120 region was then confirmed with SINFONI spectro-imaging observations (Martins et al. \\cite{mar10}). ZA07 also highlighted the porous nature of the photodissociation region (their Fig.~13), which allows the stellar radiation to leak into the ambient medium. DE09 studied this region using APEX-LABOCA 870\\,$\\mu$m and $\\it{Spitzer}$-MISPGAL observations. They refined the discussion about YSO properties by adding the 24\\,$\\mu$m MIPSGAL flux for the detected YSOs. They derived a mass of about 2000~M$_{\\odot}$ for the shell surrounding RCW~120 and proposed the existence of a Class~0 object towards the most massive condensation (their condensation~1, M$\\simeq$500~M$_{\\odot}$). However, it was not possible at that time to characterize the properties of this source, which had only been detected in the MIPSGAL 70\\,$\\mu$m data, because no data were available at longer wavelengths. The far-IR range is indeed crucial for constraining the properties of YSOs, especially by means of spectral energy distribution (SED) fitting. Another point impossible to address at that time was the deeply embedded, young stellar population that may form in high-extinction regions, which had not been detected at 24\\,$\\mu$m. A complete census of the YSOs is needed to understand how star formation proceeds in this region. In the framework of triggered star formation it is particularly important to search for an age sequence as a function of distance from the ionization front, which may identify younger stars far from it. With its unprecedented resolution and sensitivity, the $\\it{Herschel}$ Space Observatory (Pilbratt et al.~\\cite{pil10}) allows us to complete this study.  Here we present the first $\\it{Herschel}$-PACS and -SPIRE images of RCW~120. We use these images to discuss the distribution and physical properties of the YSOs observed in the field. Anderson et al.~({\\cite{and10}) discuss the dust properties in this region. ",
        "conclusions": "We have presented the first images obtained with PACS and SPIRE for the Galactic \\HII\\ region RCW~120. These images detected an embedded population of YSOs that had not been previously detected, or had been barely detected, at wavelengths shorter than 100\\,$\\mu$m. The broad wavelength coverage of PACS and SPIRE allows us to constrain the SED of these sources and derive their physical properties. Our main conclusions are: \\begin{itemize} \\item{We detect a massive (8--10~M$_{\\odot}$) YSO towards condensation 1 and show that this source is a massive Class~0 object. This is the first detection of a massive Class~0 formed by the collect and collapse process on the borders of an \\HII\\ region. This source was barely detected previously in the MIPSGAL 70\\,$\\mu$m data but its mass and exact evolutionary stage were unknown. } \\item{Our stronger constraints on the source SEDs allow us to determine more accurately the physical properties of the observed YSOs. This provides additional insight into the star formation history of this region because we can now trace (and compare) objects with similar masses but in different evolutionary stages.} \\item{YSOs revealed by the PACS and SPIRE are observed farther away from the ionization front, especially towards condensation 5 where dark filaments are observed in the near- and mid-IR. This suggests that triggered star formation may occur by means of the tunnelling of radiation from the ionizing star of RCW~120 into the ambient medium.} \\end{itemize} Six other Galactic \\HII\\ regions that show evidence of triggered star formation at their borders will be observed with PACS and SPIRE, in imaging and spectroscopy, as part of the HOBYS and ISM GT KPs. This will allow us to study in more detail the star formation history of these regions and place stronger constraints on the mechanism responsible for forming stars. On a global Galactic scale, the Hi-GAL Open Time Key Program (Molinari et al.~\\cite{mol10}) will allow us to study the efficiency of this process in our Galaxy (Zavagno et al.~\\cite{zav10})."
    },
    "1005/1005.5036_arXiv.txt": {
        "abstract": " ",
        "introduction": "The 22-m Mopra radio telescope is one element of the Australia Telescope (Frater, Brooks \\& Whiteoak 1992). The Mopra observatory is located near Coonabarabran, New South Wales, Australia. The telescope is located at a longitude of 149\\degr\\,5\\arcmin\\,58\\arcsec\\ and a latitude of 31\\degr\\,16\\arcmin\\,58\\arcsec\\ and at an elevation of 866 metres above sea level.\\footnote{See http://www.narrabri.atnf.csiro.au/mopra/ for more details.}  Construction of the wheel-on-track telescope (Cooper et al.\\ 1992) was completed in late 1991, with the first astronomical observations being made as part of a VLBI session the following week and reported in Tingay et al.\\ (1998). In its first years, Mopra was used for VLBI and single-dish observations at centimetre wavelengths (e.g., Koribalski, Johnston \\& Otrupcek 1994). In 1995, the telescope was used together with the Parkes 64-m telescope as part of Project Phoenix, which observed 209 solar-type stars over the frequency range 1.2--3.0\\,GHz to search for possible narrow-band transmissions from extra-terrestrial life (Tarter 1997). The Mopra 20/13\\,cm feed and receiver was modified from the ATCA design to improve performance over the wider Project Phoenix range --- changes which later enabled Mopra to support VLBI observations during the descent of the Huygens probe through the atmosphere of Titan (Witasse et al.\\ 2006).  A 3-mm SIS receiver was tested on the antenna in late 1994, with the feed illuminating the central solid-paneled, 15\\,m of the dish, with astronomical 3\\,mm observations beginning the following year (e.g., Elmouttie, Haynes, \\& Jones, 1997; Liang et al.\\ 1997). A 12-mm receiver was constructed for Mopra in 1996, primarily to support space VLBI observations with the HALCA satellite as part of the VLBI Space Observatory Programme (Hirabayashi et al.\\ 2000). The compromised performance of the HALCA 12-mm receiver (Kobayashi et al.\\ 2000) resulted in Mopra support for the VSOP mission being at 1.6 and 4.9\\,GHz (e.g., Scott et al. 2004; Dodson et al. 2008), however the ``VSOP'' 12-mm receiver was used in a number of other Long Baseline Array VLBI observations (e.g., Greenhill et al.\\ 2003) until it was replaced with an MMIC-based receiver (Gough et al.\\ 2004) in 2006.  Under an agreement between ATNF and the University of New South Wales (UNSW), UNSW funded the replacement of the outer perforated panels of the dish in 1999 with solid panels, and the receiver optics were modified to illuminate the full 22\\,m. Several rounds of holography using a 30-GHz satellite beacon were used to reduce the surface RMS error from 270\\,$\\mu$m in 1999 to 180\\,$\\mu$m in 2004. In 2004, observing capabilities were enhanced with the addition of an on-the-fly mapping mode, enabling efficient mapping of larger regions of interest. The 3-mm SIS receiver was the main receiver used for single-dish observing during this period, with the telescope being characterised in the 3-mm band by Ladd et al.\\ (2005). Following the replacement of the 3-mm SIS receiver in 2005, and the addition of a 12-mm receiver in 2006 and 7-mm receiver (Moorey et al.\\ 2008) in 2008, Mopra now has a suite of three receivers for single-dish observations, covering the frequency ranges 16--28, 30--50 and 76--117\\,GHz.  The receivers are cryogenically cooled Indium Phosphide (InP) High Electron Mobility Transistor (HEMT) Monolithic Microwave Integrated Circuits (MMICs) low-noise amplifiers. The receiver systems do not require any manual tuning and include noise diodes for system temperature determination, with the 3\\,mm receiver also having a paddle for atmospheric attenuation ($\\tau$) corrections. Physically, the three receivers are mounted in a single dewar (Moorey et al.\\ 2008), with a short turret rotation required to change between 12-mm and 3-mm observing frequencies. Changing to and from 7-mm observing requires a physical translation of the dewar by about 10\\,cm, which takes several minutes.  In 2006, the Mopra Spectrometer (MOPS)\\footnote{MOPS was funded in part through a grant instigated by the University of New South Wales (UNSW) under the Australian Research Council Grants scheme for Linkage, Infrastructure, Equipment and Facilities (LIEF), and in part by CSIRO ATNF.}, was installed as the main back-end for single-dish observing. MOPS is a broadband digital filter bank with four overlapping sub-bands each with a bandwidth of 2.2\\,GHz providing a total of 8.3\\,GHz of continuous bandwidth. The spectrometer can be used in two observing modes: wideband mode and zoom mode. The wideband mode provides 8.3\\,GHz of continuous coverage with a total of 32,384 spectral channels in each polarisation, providing a frequency resolution of 256\\,kHz per channel and a velocity resolution of 3.4, 1.7 and 0.9\\,\\kms\\ at 22, 44 and 90\\,GHz, respectively. The second of these modes provides a total of 16 zoom windows --- four of which can be placed in each of the four sub-bands --- each with 137\\,GHz of bandwidth and 4,096 spectral channels providing approximately eight times higher frequency and velocity resolution than the wideband mode.  Observations at 3\\,mm are restricted to the winter months (May--October) when the atmosphere is generally more stable and the amount of water vapour in the atmosphere is generally lower. However, observations at 12 and 7\\,mm are less affected by atmospheric fluctuations and water vapour content than at 3\\,mm, and thus, can be scheduled through most of the year. Therefore the installation of the 12-  and 7-mm receivers has significantly increased the potential science output of the telescope (e.g., Urquhart et al.\\ 2009, Walsh et al.\\ 2007, 2008).   In this paper we will present a combination of continuum and spectral-line observations of Jupiter and H$_2$O and SiO masers chosen to characterise the telescope's beam size, shape and efficiency between 16--50\\,GHz. To maintain consistency across all three millimetre bands we have followed the procedures and analyses used by Ladd et al.\\ (2005) to characterise the 3-mm system. ",
        "conclusions": "\\begin{tabular}{clcc} \\hline Frequency & Beam & Angular & Encircled \\\\ & Component&Size & Power\\\\ (GHz)& &(\\arcmin) & (\\%)\\\\ \\hline 22\t&\tMain\t&0.0--2.8\t&83\t\\\\ &\tInner\t&2.8--6.5\t&13\t\\\\ &\tOuter\t&6.5--9.5\t&4\t\\\\ 43\t&\tMain\t&0.0--1.3\t&77\t\\\\ &\tInner\t&1.3--3.2\t&15\t\\\\ &\tOuter\t&3.2--5.5\t&8\t\\\\   \\hline \\end{tabular} \\end{center} \\end{table}   \\begin{figure*}[!ht] \\begin{center} \\includegraphics[width=0.32\\textwidth, trim=50 0 50 0]{2009-11-18_0812-M999-M999-bandpass_4_72_17GHz.ps} \\includegraphics[width=0.32\\textwidth, trim=50 0 50 0]{2009-11-18_0812-M999-M999-bandpass_3_72_19GHz.ps} \\includegraphics[width=0.32\\textwidth, trim=50 0 50 0]{2009-11-18_0838-M999-M999-bandpass_4_72_21GHz.ps}\\\\ \\includegraphics[width=0.32\\textwidth, trim=50 0 50 0]{2009-11-18_0838-M999-M999-bandpass_3_72_23GHz.ps} \\includegraphics[width=0.32\\textwidth, trim=50 0 50 0]{2009-11-18_0838-M999-M999-bandpass_2_72_25GHz.ps} \\includegraphics[width=0.32\\textwidth, trim=50 0 50 0]{2009-11-18_0838-M999-M999-bandpass_1_72_27GHz.ps}  \\caption{Maps of the telescope beam produced from observations of Jupiter using the 12\\,mm receiver. Each map covers a 2\\,GHz frequency window and is separated by 2\\,GHz interval through the 16--28\\,GHz frequency range of the 12-mm receiver. Contours have been added to illustrate the shape of the beam at the various frequencies; these start at 10\\% of the peak emission and increase in steps of 20\\%.}\\label{fig:12mm_beam_maps} \\end{center} \\end{figure*}   Inspection of both the H$_2$O and SiO maser integrated emission map shown in Figure\\,3 (left panels) reveals a central core and two weaker azimuthal components resulting from the telescope sidelobes. Following Ladd et al.\\,(2005) we refer to these as the `main beam', `inner error beam' and `outer error beam'. All three structures can be clearly distinguished and appear to be circular in shape. The azimuthally averaged radial plots presented in the right panels of Figure\\,3 of the two maser transitions reveals similar patterns. The central `core' emission in both sets of plots are well-fit with Gaussians of FWHM 2.35\\arcmin\\ and 1.2\\arcmin\\ for the H$_2$O and SiO maser respectively. In both cases the main beam sizes determined by the fit to the maser emission is approximately the same as determined using the cross-scan data presented in Section\\,3.1. Comparing the expect Airy patterns with the data we see the main beam and, inner and outer error beam for the two masers correspond well with the expected position and intensity ratios of the first, second and fourth components of the Airy pattern. However, there is no counterpart in either sets of data that would correspond with the position of the third Airy component, and it is unclear what might be responsible for the suppression of this particular sidelobe.    Measuring the total encircled power within the main beam at 22\\,GHz and 43\\,GHz we find values of $\\sim$83\\% and $\\sim$77\\% respectively, which is marginally better than the 75\\% reported by Ladd et al. (2005) in their analysis of the telescope's 3\\,mm beam pattern. However, this means that the error beams can make a significant contribution to the measured flux for extended sources and need to be taken into account when estimating the beam efficiency (this will be discussed in more detail in Section 3.3.2). In Table\\,\\ref{tbl:beam_summary} we present a summary of the angular extent and the total power encircled in the main, and inner and outer error beams.    \\subsection{Beam Efficiency}    In Figures\\,\\ref{fig:12mm_beam_maps} and \\ref{fig:7mm_beam_maps} we present the results of our beam mapping observations of Jupiter. The emission maps nicely illustrate the circular shape and the decrease of the telescope's FWHM beam size with increasing frequency, as expected. We note that in a number of the maps the position of the planet is offset from the pointing centre (0,0), which is the result of pointing errors. These are bigger than would normally be expected due to the large angular distance of the planet from the Galactic Plane during the observations, and consequently, the large angular distances between Jupiter and the nearest available pointing sources.  As we have seen in the previous section the radial beam profile at both 12 and 7\\,mm consists of three distinct components. The presence of these error beams somewhat complicates the calibration for sources with angular diameters larger than about twice the telescope's main beam. A consequence of this is that the telescope efficiency will depend on the angular size of the observed source. Following the lead of Ladd et al. (2005), we separate the calibration into two parts to accommodate both compact and extended sources.   \\subsubsection{Main Beam Efficiency}    \\begin{figure} \\begin{center} \\includegraphics[width=0.23\\textwidth, trim=50 0 50 0]{2009-10-19_1122-M347-M999-bandpass_4_61_31GHz.ps} \\includegraphics[width=0.23\\textwidth, trim=50 0 50 0]{2009-10-19_1122-M347-M999-bandpass_3_61_33GHz.ps}\\\\ \\includegraphics[width=0.23\\textwidth, trim=50 0 50 0]{2009-10-14_1020-M347-M999-bandpass_4_72_35GHz.ps} \\includegraphics[width=0.23\\textwidth, trim=50 0 50 0]{2009-10-14_1020-M347-M999-bandpass_3_72_37GHz.ps}\\\\ \\includegraphics[width=0.23\\textwidth, trim=50 0 50 0]{2009-10-14_1020-M347-M999-bandpass_2_72_39GHz.ps} \\includegraphics[width=0.23\\textwidth, trim=50 0 50 0]{2009-10-14_1020-M347-M999-bandpass_1_72_41GHz.ps}\\\\ \\includegraphics[width=0.23\\textwidth, trim=50 0 50 0]{2009-10-14_1056-M347-M999-bandpass_1_66_43GHz.ps} \\includegraphics[width=0.23\\textwidth, trim=50 0 50 0]{2009-10-14_1056-M347-M999-bandpass_2_66_45GHz.ps}\\\\ \\includegraphics[width=0.23\\textwidth, trim=50 0 50 0]{2009-10-14_1056-M347-M999-bandpass_3_66_47GHz.ps} \\includegraphics[width=0.23\\textwidth, trim=50 0 50 0]{2009-10-14_1056-M347-M999-bandpass_4_66_49GHz.ps}  \\caption{Maps of the telescope beam produced from observations of Jupiter using the 7-mm receiver. Each map covers a 2\\,GHz frequency window and are separated by 2\\,GHz interval through the 30--50\\,GHz frequency range of the 7-mm receiver. Contours have been added to illustrate the shape of the beam at the various frequencies; these start at 10\\% of the peak emission and increase in steps of 20\\%.}\\label{fig:7mm_beam_maps} \\end{center} \\end{figure}    Planets are ideal astronomical sources to use for deriving beam efficiencies as their sizes and fluxes are well known. Moreover, given that Jupiter is unresolved at all frequencies of interest for this study it couples well with the telescope's main beam but not to either of the two error beams. In this cases the measured antenna temperature, $T^*_{\\rm{A}}$, is a convolution of the sources brightness temperature distribution, $T_{\\rm{B}}(\\theta,\\phi)$, with the telescope beam pattern, $P_{\\rm{n}}(\\theta,\\phi)$, such that:  \\begin{equation} T^*_{\\rm{A}}=\\frac{\\int T_{\\rm{B}}(\\theta,\\phi) P_{\\rm{n}}(\\theta,\\phi)\\,{\\rm{d}}\\Omega}{\\int_{\\rm{4\\pi}} P_{\\rm{n}}(\\theta,\\phi) \\,{\\rm{d}}\\Omega} \\end{equation}   \\begin{table*}[!ht] \\begin{center} \\caption{Summary of measured antenna temperatures, correction factors and beam efficiencies as a function of frequency. The columns are as follows: (1) observing frequency, (2) measured peak antenna temperatures, (3) and (4) beam dilution and coupling factors , (5) planet's diameter, (6) two numbers are presented, the first of which refers to the planets physical temperature taken from the model of de Pater \\& Massie\\,(1985), while the second number refers to the  Rayleigh-Jeans correction that needs to be applied to obtain the planets brightness temperature at a given frequency (i.e., $T_{\\rm{B}}=T_{\\rm{phys}}-\\frac{h\\nu}{2k}$), (7) expected antenna temperature for a prefect telescope, (8) conversion factor for antenna temperature and the flux density scale, (9) beam efficiency, $\\eta_{\\rm{mb}}$, calculated for compact sources which couple only to the telescopes main beam and  (10) the extended beam efficiency, $\\eta_{\\rm{xb}}$, for more extended sources ($>$ twice the FWHM of the main beam). The quoted errors are derived from the Gaussian fit to the azimuthal emission profile, however, taking into account systematic errors and atmospheric attenuation we estimate a more realistic error to be of order 10\\%.}\\label{tbl:efficiencies} \\begin{tabular}{cccccccccc} \\hline & \t &  &  & \\multicolumn{2}{c}{Planet Properties} & &Conversion\\\\ \\cline{5-6} Freq. \t& $T^*_{\\rm{A}}$\t & $f$  & $K$ & D   & $T_{\\rm{B}}$        & $T^{*\\prime}_{\\rm{A}}$&  Factor & $\\eta_{\\rm{mb}}$& $\\eta_{\\rm{xb}}$\\\\ (GHz)\t\t& (K)\t\t\t\t &     &  &(\\arcsec)  & (K)&    (K)     &(Jy K$^{-1}$)                   &                 \\\\ \\hline (1)\t& (2)\t & (3)  & (4) & (5)   & (6)        & (7)&  (8) & (9)& (10) \\\\ \\hline \\hline  17\t&\t3.16$\\pm$0.12\t&\t22.3\t&\t1.023\t&\t39.5\t&\t155$-$0.41\t&\t6.77\t&\t14.1\t&\t0.47$\\pm$0.04\t&\t0.56$\\pm$0.05\t\\\\ 19\t&\t3.85$\\pm$0.14\t&\t18.9\t&\t1.027\t&\t39.5\t&\t146$-$0.46\t&\t7.50\t&\t12.9\t&\t0.51$\\pm$0.04\t&\t0.62$\\pm$0.04\t\\\\ 21\t&\t4.24$\\pm$0.14\t&\t16.5\t&\t1.031\t&\t39.5\t&\t138$-$0.50\t&\t8.07\t&\t12.8\t&\t0.53$\\pm$0.03\t&\t0.63$\\pm$0.04\t\\\\ 23\t&\t5.00$\\pm$0.16\t&\t14.0\t&\t1.036\t&\t39.5\t&\t134$-$0.55\t&\t9.17\t&\t12.3\t&\t0.54$\\pm$0.03\t&\t0.65$\\pm$0.04\t\\\\ 25\t&\t6.19$\\pm$0.16\t&\t12.9\t&\t1.039\t&\t39.5\t&\t135$-$0.60\t&\t10.00\t&\t10.9\t&\t0.62$\\pm$0.03\t&\t0.74$\\pm$0.03\t\\\\ 27\t&\t7.77$\\pm$0.18\t&\t10.3\t&\t1.049\t&\t39.5\t&\t138$-$0.65\t&\t12.74\t&\t11.1\t&\t0.61$\\pm$0.02\t&\t0.73$\\pm$0.03\t\\\\ \\hline 31\t&\t12.05$\\pm$0.18\t&\t5.3\t&\t1.097\t&\t43.1\t&\t146$-$0.74\t&\t24.80\t&\t14.4\t&\t0.49$\\pm$0.02\t&\t0.63$\\pm$0.02\t\\\\ 33\t&\t14.55$\\pm$0.15\t&\t4.7\t&\t1.110\t&\t43.1\t&\t149$-$0.79\t&\t28.32\t&\t13.7\t&\t0.51$\\pm$0.01\t&\t0.67$\\pm$0.01\t\\\\ 35\t&\t16.26$\\pm$0.18\t&\t4.4\t&\t1.119\t&\t43.5\t&\t152$-$0.84\t&\t30.97\t&\t13.6\t&\t0.52$\\pm$0.01\t&\t0.68$\\pm$0.01\t\\\\ 37\t&\t17.87$\\pm$0.15\t&\t4.0\t&\t1.129\t&\t43.5\t&\t154$-$0.89\t&\t33.60\t&\t13.6\t&\t0.53$\\pm$0.01\t&\t0.69$\\pm$0.01\t\\\\ 39\t&\t18.68$\\pm$0.16\t&\t3.7\t&\t1.141\t&\t43.8\t&\t156$-$0.94\t&\t36.75\t&\t14.4\t&\t0.51$\\pm$0.01\t&\t0.66$\\pm$0.01\t\\\\ 41\t&\t21.21$\\pm$0.17\t&\t3.3\t&\t1.161\t&\t43.8\t&\t158$-$0.98\t&\t41.41\t&\t14.5\t&\t0.51$\\pm$0.01\t&\t0.67$\\pm$0.01\t\\\\ 43\t&\t21.47$\\pm$0.17\t&\t3.4\t&\t1.154\t&\t43.1\t&\t160$-$1.03\t&\t40.51\t&\t14.1\t&\t0.53$\\pm$0.01\t&\t0.69$\\pm$0.01\t\\\\ 45\t&\t22.27$\\pm$0.17\t&\t3.1\t&\t1.167\t&\t43.1\t&\t161$-$1.08\t&\t43.52\t&\t14.8\t&\t0.51$\\pm$0.01\t&\t0.67$\\pm$0.01\t\\\\ 47\t&\t23.10$\\pm$0.24\t&\t3.0\t&\t1.177\t&\t43.1\t&\t163$-$1.13\t&\t46.02\t&\t15.3\t&\t0.50$\\pm$0.01\t&\t0.65$\\pm$0.01\t\\\\ 49\t&\t21.61$\\pm$0.19\t&\t2.7\t&\t1.196\t&\t43.1\t&\t164$-$1.18\t&\t50.27\t&\t18.1\t&\t0.43$\\pm$0.01\t&\t0.56$\\pm$0.01\t\\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  For compact sources with a uniform brightness temperature the equation above simplifies to:  \\begin{equation} T^*_{\\rm{A}} = \\eta_{\\rm{mb}} \\frac{\\int_{\\rm{planet}} P_{\\rm{n}}(\\theta,\\phi)\\,{\\rm{d}}\\Omega}{\\int_{\\rm{mb}} P_{\\rm{n}}(\\theta,\\phi)\\,{\\rm{d}}\\Omega} T_{\\rm{B}} = \\eta_{\\rm{mb}} \\frac{\\Omega_{\\rm{planet}}} {\\Omega_{\\rm{mb}}} T_{\\rm{B}} \\end{equation}  \\noindent where $\\Omega_{\\rm{planet}}$ is the solid angle of the planet at the time of the observations, and $\\Omega_{\\rm{mb}}$ is the telescope main beam solid angle. $\\eta_{\\rm{mb}}$ is referred to as the main beam efficiency and is defined as the ratio of the total measured power within the telescope's main beam, and total measured power integrated over the full sphere of 4$\\pi$, i.e.:   \\begin{equation} \\eta_{\\rm{mb}} = \\frac{\\int_{\\rm{mb}} P_{\\rm{n}}(\\theta,\\phi)\\,{\\rm{d}}\\Omega}{\\int_{4\\pi} P_{\\rm{n}}(\\theta,\\phi)\\,{\\rm{d}}\\Omega} % \\end{equation}  The ratio of planetary and main beam solid angles shown in the right hand side of Equation\\,3 consists of two important correction factors: beam dilution, $f$, which reduces the expected antenna temperature --- this can be large if the angular size of the main beam is much larger than angular size of the planet; and the coupling coefficient, $K$, which corrects for the antenna's response for the non-pointlike planetary disk. Beam dilution is simply the ratio of the area of the planetary disk ($\\pi/4\\times D^2_{\\rm{planet}}$ --- where $D$ is the angular diameter of the planet) and the telescope's main beam ($1.133\\times \\theta^2_{\\rm{mb}}$):  \\begin{equation} f^{-1}=\\frac{D^2_{\\rm{planet}}} {\\theta^2_{\\rm{mb}}}{\\rm{ln}} 2 \\end{equation}  \\noindent and $K$ can be calculated following Baars (1973) Equation\\,12:  \\begin{equation} K = \\frac{x^2}{1-{\\rm{exp}}(-x^2)}\\,{\\rm{where}}\\,x=\\frac{D_{\\rm{planet}}}{\\theta_{\\rm{mb}}}\\sqrt{{\\rm{ln}}2} \\end{equation}  Both correction factors $K$ and $f$ have been derived using the angular diameter of the planet at the time of the observation, and the  FWHM of the telescopes main beam has been derived from a Gaussian fit to the azimuthal radial profiles of the Jupiter maps (i.e., using Equation\\,1). The derived values of $K$ and $f$ and the angular size of the planet at the time of the various observations are presented in Columns\\,3, 4 and 5 of Table\\,3, respectively.  \\begin{figure*} \\begin{center} \\includegraphics[width=0.95\\textwidth]{efficiency_plot_correct_deconvolved.eps}  \\caption{Measurements of the Mopra telescope beam efficiencies for the 12 and 7-mm receivers. The vertical error bars result from the propagation of the 1$\\sigma$ error associated with the measured peak antenna temperature, while the horizontal error bars simply indicate the frequency bandpass over which the measurement is applicable.  The regions shaded in grey indicate the frequency range covered by the 12- and 7-mm receivers and are labelled accordingly. The results of a polynomial fit to the measured main beam efficiency ($\\eta_{\\rm{mb}}$) for each frequency range are plotted in red. The blue line indicates the estimated efficiency of the extended beam (see Section 3.3.2 for details).}\\label{fig:efficiency_plot} \\end{center} \\end{figure*}  \\begin{sloppypar} The frequency-dependent brightness temperature of Jupiter is derived from the model of de Pater \\& Massie\\,(1985), which gives the physical temperature, $T_{\\rm{phys}}$.  The planet's brightness temperature is obtained by applying the Rayleigh-Jeans correction: $T_{\\rm{B}}=T_{\\rm{phys}}-\\frac{h\\nu}{2k}$ --- the planets physical temperature and Rayleigh-Jeans correction are given in this form in Column\\,6 in Table\\,3. The correction is $\\sim$0.5 and 1\\,K for the 12- and 7-mm receiver ranges respectively. This model has been adopted because it yields brightness temperatures at all frequencies and is also used for calibration at the Australia Telescope Compact Array. The model is quite consistent with recent broadband measurements of the brightness temperature of Jupiter: 152$\\pm$5\\,K at 32\\,GHz (6\\,GHz bandwidth, Mason et al.\\,1999); 134$\\pm$4, 146.6$\\pm$2.0, and 154.7$\\pm$1.7\\,K at 23, 33, and 41\\,GHz, respectively (Page et al.\\ 2003). Jupiter's brightness temperature, the beam dilution and coupling coefficient together yield the expected antenna temperature, ${T^{*\\prime}_{\\rm{A}}}$, assuming a perfect dish and loss-less front and backends, i.e.; \\end{sloppypar}  \\begin{equation} T^{*\\prime}_{\\rm{A}} = \\frac{T_{\\rm{B}}}{fK} \\label{eqn:efficiency} \\end{equation}  \\begin{sloppypar} The expected antenna temperatures calculated using the equation above are presented in Column\\,7 of Table\\,3.  In Column\\,8 of Table\\,3 we present the factor required to convert between antenna temperature and the flux density scale for point source estimated from our Jupiter measurements. Finally, we estimate the main beam efficiency using: \\end{sloppypar}  \\begin{equation} \\eta_{\\rm{mb}}= \\frac{T^*_{\\rm{A}}}{T^{*\\prime}_{\\rm{A}}} \\label{eqn:efficiency} \\end{equation}  \\noindent The main beam efficiencies can be found in Column\\,9 of Table\\,3.  In Figure\\,\\ref{fig:efficiency_plot} we present a plot of the main beam efficiency as a function of frequency. The individual data points are shown as filled circles and the frequency ranges for each of the receivers are indicated by grey shading. The main beam efficiencies obtained for the 12-mm receiver show an increasing efficiency with frequency beginning at slightly below 50\\% at 16\\,GHz rising to $\\sim$65\\% at the upper frequency range of 27\\,GHz. The 12-mm data can be reasonably well fit with a straight line $y=A+Bx$ where $A=21.76$ and $B=1.49$ and $x$ is the observed frequency in GHz.  The profile of the beam efficiencies over the 7-mm range is a little more complicated. It appears reasonably flat between 33 and 47\\,GHz with an average main beam efficiency of $\\sim$50\\%, with the efficiency falling off towards either end of the frequency band. The 7-mm data is best fit with a fourth order polynomial, $y=A+Bx+Cx^2+Dx^3+Ex^4$ where $A=-4085.61$,  $B=425.97$, $C=-16.3656$, $D=0.27813$ and $E=-0.00176445$ and $x$ is as previously defined.   \\subsubsection{Mopra Extended Beam Efficiency}   As we have seen in Section\\,3.2 the telescope beam pattern consists of three components we referred to as the main beam, and inner and outer error beams. Compact sources couple to the main beam only. However, extended sources (larger than $\\sim$2.8\\arcmin\\ and 1.6\\arcmin\\ at 12- and 7-mm respectively) can couple to the inner and outer error beams as well, making the calibration more complicated. Following the method of Ladd et al.\\,(2005) we write an expression for the measured antenna temperature for extended sources as a function of the extended beam pattern, $P_{\\rm{xb}}$, which included contribution from both the main beam and error beams such that:  \\begin{equation} T^*_{\\rm{A}} = \\eta_{\\rm{xb}} \\frac{\\int T_{\\rm{B}}(\\theta,\\phi) P_{\\rm{xb}}(\\theta,\\phi)\\,{\\rm{d}}\\Omega}{\\int_{\\rm{mb}} P_{\\rm{xb}}(\\theta,\\phi)\\,{\\rm{d}}\\Omega} \\end{equation}  \\noindent where  $\\eta_{\\rm{xb}}$ is the extended beam efficiency defined as:  \\begin{equation} \\eta_{\\rm{xb}} = \\frac{\\int_{\\rm{xb}} P_{\\rm{xb}}(\\theta,\\phi)\\,{\\rm{d}}\\Omega}{\\int_{4\\pi} P_{\\rm{n}}(\\theta,\\phi)\\,{\\rm{d}}\\Omega} % \\end{equation}  Similarly to Ladd et al. (2005) we have no direct measurements of extended sources, however, we can use the radial beam patterns obtained from the maser beam maps presented in Figure\\,3 to estimate the extended beam efficiency. Referring to Table\\,2, we estimate the amount of power within the inner and outer error beams is approximately 20\\% and 24\\% of that found in the main beam at 22 and 43\\,GHz, respectively. These values would be reduced to 15\\% at both frequencies if the source couples to the main beam and the inner error beam only. Using our knowledge of the radial beam pattern we estimate the fraction of the total power contained within the extended beam is $\\sim$60--75\\% for the 12-mm receiver range and $\\sim$65\\% for the 7-mm receiver. We present the values of the extended beam efficiency in the final column of Table\\,3 and are also shown on Figure\\,6 (blue line)."
    },
    "1005/1005.0425_arXiv.txt": {
        "abstract": "FBS0107-082 is an emission line object previously classified as a nova-like cataclysmic variable star.  New optical spectroscopy shows very strong hydrogen Balmer lines, along with a nebular forbidden line spectrum and absorption features from an early-F photosphere.  When combined with other IR and optical data from the literature, these data point to the object being a symbiotic nova seen in a prolonged outburst.  Photometry on time scales of minutes, days, and years show only very weak variability. ",
        "introduction": "FBS~0107-082 (also known as Cet 7, and designated FBS0107 in this paper) was found in the First Byurakan Survey (FBS) for objects with UV excesses.  Based on spectra obtained with the 6-m telescope at the Special Astronomical Observatory, Kopylov et al. (1988) provided a preliminary classification as a nova-like (NL) cataclysmic variable (CV).  FBS0107 was subsequently listed as a CV in SIMBAD and in the Downes, Webbink, \\& Shara (1997) CV catalog, but has received very little attention since the initial discovery paper.  We obtained spectra of FBS0107 in 2005 as part of a survey for winds in NL CVs, finding that the spectrum resembles that of a symbiotic star or planetary nebulae.  We have complemented these initial spectra with multi-epoch photometry and with an analysis of archival ESO spectra.  In this paper we report on our spectroscopic and photometric findings and discuss the most likely kinds of objects that fit these findings. The star appears to be an unusual symbiotic binary.  Symbiotic stars (Kenyon 1986) have evidence for a hot component in the near UV plus a cool component in the near-IR. In most cases the hot component is judged to be a white dwarf or evolved subdwarf, and the companion to be an M giant with a strong wind.  Nebular emission lines are produced by the action of the radiation field of the hot component on the surrounding plasma from the wind.  Symbiotic stars are not a particularly homogeneous group and the evolutionary state of the hot component (e.g. subdwarf or white dwarf) is often uncertain.  Our observations are not a perfect fit to the symbiotic classification, but this identification seems to best match the properties of FBS0107. ",
        "conclusions": "Although our conclusion is not as rugged as desired, our opinion is that FBS0107 is most likely a symbiotic nova, seen in prolonged outburst.  This judgement is partly due to the shortcomings of various competing identifications, but is mostly based on the presence of the early-F photosphere in the spectrum of FBS0107.  The presence of this absorption spectrum is quite secure and is difficult to explain without invoking the outburst phase of a nova event.  The fact that FBS0107 deviates somewhat from the properties of most symbiotic stars insofar as IR colors and emission line ratios is tentatively attributed to the fact that FBS0107 is in prolonged outburst, and therefore not representative of the majority of symbiotic binaries which are found in quiesence. A somewhat worrisome implication of our interpretation is that FBS0107 must be an extreme member of the small class of symbiotic novae, having an outburst time scale perhaps as long as that of AG Peg, and changing so slowly at present that photometric variations are barely detectable.  Symbiotic binaries have been considered as progenitors of Type Ia supernovae (e.g. Munari \\& Renzini 1982; Kenyon et al. 1993; Hachisu et al. 1999), especially the subclasses of recurrent novae (Mikolajewska 2008) and symbiotic novae, (e.g. Hernanz \\& Jos\\'{e} 2008).  Note however that there are considered  arguments against symbiotic binaries being a significant channel for SNIa (e.g. Hillebrandt \\& Niemeyer 2000; Iben \\& Fujimoto 2008).  In any case, this aspect of symbiotic binaries may lend additional motivation to understanding the nature of FBS0107, especially if FBS0107 is able to remain in very protracted thermonuclear outburst, as we are suggesting.  Fortunately our symbiotic outburst hypothesis has fairly straightforward tests.  If FBS0107 is currently in symbiotic outburst, then photographic plate archives may provide a record of the pre-outburst state.   Also, narrow-band imaging at high angular resolution may reveal nebulosity. Non-optical observations can further discriminate among the possible candidates for the nature of FBS0107; for example, it would be most interesting to know if FBS0107 is (or has been) a supersoft x-ray source, indicative of steady nuclear burning.  Overall, it seems that one of our more definite conclusions is the all-too-common ending to observational papers that ``more observations are needed''."
    },
    "1005/1005.0563_arXiv.txt": {
        "abstract": "This paper makes two points. First, we show that the line-of-sight solution to cosmic microwave anisotropies in Fourier space, even though formally defined for arbitrarily large wavelengths, leads to position-space solutions which only depend on the sources of anisotropies inside the past light-cone of the observer. This foretold manifestation of causality in position (real) space happens order by order in a series expansion in powers of the visibility $\\gamma = e^{-\\mu}$, where $\\mu$ is the optical depth to Thompson scattering. We show that the contributions of order $\\gamma^N$ to the CMB anisotropies are regulated by spacetime window functions which have support only inside the past light-cone of the point of observation. Second, we show that the Fourier-Bessel expansion of the physical fields (including the temperature and polarization momenta) is an alternative to the usual Fourier basis as a framework to compute the anisotropies. The viability of the Fourier-Bessel series for treating the CMB is a consequence of the fact that the visibility function becomes exponentially small at redshifts $z \\gg 10^3$, effectively cutting off the past light-cone and introducing a finite radius inside which initial conditions can affect physical observables measured at our position $\\vec{x}=0$ and time $t_0$. Hence, for each multipole $\\ell$ there is a discrete tower of momenta $k_{i \\, \\ell}$ (not a continuum) which can affect physical observables, with the smallest momenta being $k_{1 \\, \\ell} \\sim \\ell$. The Fourier-Bessel modes take into account precisely the information from the sources of anisotropies that propagates from the initial value surface to the point of observation -- no more, no less. We also show that the physical observables (the temperature and polarization maps), and hence the angular power spectra, are unaffected by that choice of basis. This implies that the Fourier-Bessel expansion is the optimal scheme with which one can compute CMB anisotropies. ",
        "introduction": "The cosmic microwave background (CMB) is the earliest, cleanest observation that reveals what the Universe looked like at the very beginning. A remarkable string of observations of the CMB temperature fluctuations over the last 20 years, most notably by COBE-DMR \\cite{Smoot:1992td} and WMAP \\cite{Bennett:2003bz,Spergel:2003cb,Komatsu:2010fb}, has shown that the typical initial conditions of the Universe when it was under $400.000$ years old can be characterized by extreme homogeneity and isotropy, only slightly perturbed by small, ${\\cal{O}}(10^{-5})$ fluctuations with a nearly scale-invariant spectrum. More recently, the small degree of polarization that is imprinted on the CMB radiation by anisotropic Thompson scattering has also started to become detectable \\cite{Page:2006hz,Brown:2009uy,Chiang:2009xsa}, and may hold the key to unravel the mistery of the birth of our Universe -- see, e.g., \\cite{Dodelson:03,Mukhanov:05,Durrer:2008,Peter:2009zzc}.  The remarkable success of the CMB as probably the most powerful tool in observational Cosmology can also be traced to the simplicity of its underlying mechanisms: Thompson scattering and linear perturbation theory. The basic theory, which is an application of the relativistic radiative transfer equations \\cite{Chandrasekhar}, was initially developed in connection with the CMB by Peebles and Yu \\cite{Peebles:1970ag}, and the first to write down the full collisional Boltzmann equations for temperature and polarization were Bond and Efstathiou \\cite{Bond:1984fp,Bond:1987ub}.  However, even if the main mechanisms driving acoustic oscillations in the baryon-photon fluid were basically understood early on, crucial features such as neutrinos, gravitational waves, spatial curvature and the effects of lensing on polarization remained puzzling. It was not until the 1990's that the theory reached full maturity, when the complete separation of radial and angular modes allowed by the use of spin angular momentum eigenfunctions for polarization cleared the way for our current understanding of CMB physics \\cite{Ma:1995ey,Hu:1995en,Hu:1996qs,Seljak:1996gy,Zaldarriaga:1996xe,Kamionkowski:1996zd,Kamionkowski:1996ks,Hu:1997hp,Hu:1997mn}.  One particularly significant step forward was achieved with the line-of-sight solution to the collisional Boltzmann equations \\cite{Seljak:1996is}. The idea is that photons travel along null geodesics, hence the comoving distance $\\Delta x$ between two successive collisions is equal to the conformal time interval $\\Delta \\eta$ between those collisions (as usual, we assume that the Born approximation is valid.) This means that, given a line-of-sight $\\hat{n}$, we know that a photon detected at time $\\eta$ travelling along that line-of-sight was at the position $\\vec{x}'=\\hat{n} (\\eta-\\eta') $ at time $\\eta'$, if no collisions occurred between those times. The more general case of an ensemble of photons can be easily accomodated in this picture, since the probability that a photon scatters with free electrons is given in terms of the optical depth for Thompson scattering, which to a very good approximation is a smooth function of time, $\\mu(\\eta)$. The final state of the ensemble is then obtained through integration over time of the sources of temperature and polarization anisotropies, appropriately weighted by the optical depth at each time.   The power of the line-of-sight integral solution is that it separates, as much as it is possible, the (free) propagation of photons from the ultimate sources of anisotropies (matter and metric perturbations) -- so, it is similar in spirit to a Green's function for the temperature and polarization of an ensemble of photons. However, there is one feature of the generation of anisotropies which makes it impossible to completely separate the sources and the anisotropies: anisotropic Thompson scattering is itself a source of both temperature and polarization, so the process is, in some sense, non-local.   When photons scatter at a given place and time, the fluctuations in temperature and polarization that are generated as a consequence of those scatterings depend also on the quadrupole of the temperature and the polarization of the photons that were incident at that place and time. Therefore, the fact that those incident photons typically propagated to the location of the scattering from far away implies that the process is non-local -- hence the line-of-sight solution is actually a set of integral equations, at least for the lowest multipoles ($\\ell \\leq 2$). The higher multipoles ($\\ell \\geq 3$), however, can be completely determined from the lowest multipoles, which makes the line-of-sight solution a vastly superior method compared to the usual hierarchy of Boltzmann equations. Of course, for the low multipoles the integral equations are impractical, and the preferred method to compute them is to revert back to the hierarchy of Boltzmann equations, which is then truncated at a relatively low multipole that is sufficient to accurately compute the multipoles $\\ell \\leq 2$.   Nevertheless, despite the fact that the generation of anisotropies is a non-local mechanism, it is still completely causal: photons propagate along light cones between scatterings (which are basically instantaneous within the cosmological timescales.) The main drive behind this work is to clarify how causality and non-locality are manifested in the generation of the CMB. We are interested, in particular, in describing the generation of the CMB in position (real) space, since Fourier space can sometimes obfuscate the causal nature of the physical mechanisms.  We will show that, in position space, the non-local anisotropies can be resolved into explicitly causal pieces by iterating the integral line-of-sight equations. The smallness parameter in this perturbative expansion (which is reminiscent of the Dyson series of Quantum Electrodynamics \\cite{Sakurai}) is the {\\it total} visibility $\\gamma=e^{-\\mu}$, where $\\mu$ is the optical depth to Thompson scattering. The first term in that series corresponds to the last scattering of the photons before they were detected; the second term corresponds to the last two scatterings before detection; and so on.   In this series over the number of scatterings, causality is manifested at each order in terms of radial integrals corresponding to spacetime window functions which are only non-zero inside the past light-cone (PLC) of the observation point. The spacetime dependence of the $\\gamma^N$-order term is given in terms of integrals of $N+1$ spherical Bessel functions, which we demonstrate to have support only inside the PLC. The first term of the series, ${\\cal{O}}(\\gamma^1)$, corresponds simply to the Sachs-Wolfe (SW), Doppler and Integrated Sachs-Wolfe (ISW) contributions to the temperature fluctuations from the last scattering \\cite{1967ApJ...147...73S}, and evidently it is only nonzero at the surface of the light cone. The second-order terms for temperature and polarization carry the memory of the last two scatterings, and can be non-zero anywhere inside the volume of the PLC (not only on its surface). Fig. \\ref{LightCones} illustrates the structure of the light cones for one, two and three scatterings. Hence, we have shown that the line-of-sight integrals are actually {\\it retarded} Green's functions for the temperature and polarization, and we found analytical expressions for them in position space.   \\begin{figure}[t] \\includegraphics[scale=0.75]{LightCones}  \\caption{Past light-cones for photons experiencing one, two and three scatterings since the time of decoupling from matter (denoted as the base of the cone.) The terms corresponding to these interactions are respectively of order $\\gamma^1$, $\\gamma^2$ and $\\gamma^3$, where $\\gamma=e^{-\\mu}$ is the visibility.}  \\label{LightCones} \\end{figure}   The second result of this paper follows from this causal structure. The expansion in successive scatterings (or, equivalently, in orders of the visibility function) makes it clear that what matters for the CMB observables are not simply the source fields and the anisotropies as functions of position and time, $f(\\vec{x},\\eta)$, but the fields multiplied by powers of the visibility, $\\gamma(\\eta)^N f(\\vec{x},\\eta)$. Since the visibility vanishes for very early times (say, $z_i >> 10^3$), for all practical purposes that have to do with the CMB, our PLC is cut-off at $\\eta_i=\\eta(z_i)$, and all the sources are effectively zero outside the radius corresponding to that time (which in a typical flat $\\Lambda$CDM model is approximately $r_i \\approx 5 H_0^{-1}$.)   The fact that the physical fields are exponentially suppressed at the boundaries of the PLC means that the best basis for expanding the fields, as well as the anisotropies, is not the Fourier basis, which is most convenient for plane waves in $\\mathbb{R}^3$, but the Fourier-Bessel basis \\cite{Watson}, which expands functions $f(\\vec{x})$ inside a 3D sphere of finite radius $R$ into spherical harmonics $Y_{\\ell m} (\\hat{x})$ and spherical Bessel functions $j_\\ell(k_{i \\ell} x)$, with $k_{i \\ell}R$ being the $i^{th}$ root of $j_\\ell$. Both sets of basis functions are eigenvectors of the Laplacian operator in flat space, with eigenvalues ${-\\vec{k}}^2$ or $-k_{i \\ell}^2$, but they differ most significantly in that the Fourier-Bessel series establishes a discrete tower of momenta for each multipole $\\ell$, so that the smallest momentum mode is $k_{1 \\ell} \\sim \\ell/r_i$. Therefore, there is a clearly defined minimal mode that can contribute to CMB observables, and the modes above it are all discretized. The drawback is that, if the underlying fields are Gaussian in nature, the Fourier-Bessel modes, as opposed to the Fourier modes, are not statistically independent, so their covariance matrix is not diagonal. However, dynamics is one thing, statistics is another, and we will show how these issues can be separated so that we can easily recover angular power spectra which have precisely the same statistical properties of the power spectra computed in Fourier space.   The first part of this work takes an approach which is similar to that used to make constrained simulations of CMB temperature and polarization maps by Liguori, Matarrese and Moscardini \\cite{Liguori:2003mb} and by Komatsu, Spergel and Wandelt \\cite{Komatsu:2003iq} -- see also \\cite{Cornish:2003db,Liguori:2007sj,Elsner:2009md}. However, while those simulations make use of several approximation schemes in order to produce numerically viable codes, our expressions are exact and analytical. In particular, our results show that the transfer functions of \\cite{Liguori:2003mb,Komatsu:2003iq} have an invariant (or geometrical) piece which is described by our spacetime window functions, so that the transfer functions can be obtained by integration of these window functions over time with some visibility function. Another work close in spirit to ours was done by Bashinsky and Bertshinger \\cite{Bashinsky:2000uh,Bashinsky:2002vx}, who calculated the Green's function for the evolution of linear cosmological perturbations in position space but did not compute the Green's functions for the anisotropies.    A quick note on our conventions: the Fourier transform of a function $f(\\vec{x})$ is $f(\\vec{k}) = (2\\pi)^{-3/2}\\int d^3 x\\, e^{-i \\vec{k} \\cdot \\vec{x}} f(\\vec{x})$; the perturbed Friedman-Robertson-Walker metric is $ds^2 = a^2(\\eta) [-(1+2\\Phi)d\\eta^2 + (1-2\\Psi) d\\vec{x}^2 ]$; and the relationship between the rotation matrices and the spin spherical harmonics is $D^\\ell_{m, -s} (\\alpha,\\beta,\\gamma \\rightarrow 0) = \\sqrt{4\\pi/(2\\ell+1)} \\, \\, {}_{s} Y_{\\ell m}^* (\\beta,\\alpha)$, so that ${}_s Y^*_{\\ell m} = (-1)^{m-s} \\, {}_{-s} Y_{\\ell, -m}$. Finally, everything we will say in this paper concerns scalar (spin-0) density perturbations -- we will consider lensing and gravitational waves in future work. ",
        "conclusions": "In this paper we have shown how causality constraints in position space regulate which sources of anisotropies (the matter and metric perturbations) can contribute to the CMB. This causal structure is manifested order by order in a series of terms corresponding to the number of interactions that photons experienced over the past light-cone of the observer -- or, equivalently, a power series on the visibility $\\gamma = e^{-\\mu}$.   When expressed in position space, the line-of-sight integrals acquire an intuitive interpretation in terms of scatterings over the light-cones of the successive scatterings, all the way from the sources to the point and time of observation. In particular, we find that, in position space, only the sources of anisotropy that are inside our past light-cone are taken into account. This statement is exact to all orders -- as it should be, since the causal nature of the propagation of photons is the key ingredient in the line-of-sight integrals from which we started.   At each order in the power series on the visibility, the sources are weighted by spacetime window functions. These window functions can be complicated for a high number of scatterings, but they all obey a very simple rule: they vanish indentically unless some extremely simple sets of inequalities are satisfied. These inequalities have a simple geometrical interpretation: if the position of the source and the radii of the light-cones of the interactions cannot form a flat polygon, the spacetime window functions vanish. An interesting question which we did not have time to address is at what number of scatterings prior to free streaming we can cut off this series so that the error in the temperature distribution is, say, of order 1\\%.   One of the implications of these causality constraints is that, whatever the properties of the Universe outside a limiting radius $R$, the source fields do not propagate to the CMB observables -- and the line-of-sight integrals both in Fourier space and in the Fourier-Bessel expansion retain this property. In practice, this means that we can use the Fourier-Bessel framework to compute the CMB -- even though it puts the Universe in a ``spherical box'', and discards all the information outside of that box.   In the Fourier-Bessel basis, the fields are decomposed in spherical harmonics and a series of discrete eingenmodes $k_{i\\ell}$ (as opposed to the continous modes of Fourier space.) The first eigenmode for each multipole $\\ell$ is $k_{1\\ell} \\sim \\ell/R$. CMB observables are exactly the same as in the Fourier basis -- but the Fourier-Bessel basis is optimal, in the sense that it does not keep track of irrelevant modes, only the ones that contribute constructively to the physical observables.   The previous discussion implies that our results and methods are suitable for analytical and numerical studies of CMB temperature and polarization maps in models with large-scale inhomogeneities, statistical anisotropy or non-gaussianities of any kind kind \\cite{Liguori:2003mb,Komatsu:2003iq,Liguori:2007sj,Elsner:2009md}. Our results may be useful also in simulations of the CMB in the presence of topological defects. It is not clear whether the methods described here can be employed to study models with non-trivial topology \\cite{Cornish:2003db}, since in those cases the Fourier or Fourier-Bessel basis functions may not be eigenvectors of the Laplace-Beltrami operator.    \\subsection*{Acknowledgements}  The authors would like to thank J. C. A. Barata for many conversations on the properties of Bessel functions, and to Mathias Zaldarriaga for useful comments. This work was supported by FAPESP and CNPq."
    },
    "1005/1005.4554.txt": {
        "abstract": "{The burst-only source \\src\\ was discovered in a high X-ray-active state, while during an \\textsl{INTEGRAL} observations of the Galactic bulge on 2010 April 10. Pulsations at 518\\,Hz were discovered in the \\textsl{RXTE} data, confirming previous suggestions of possible associations between burst-only sources and accreting millisecond X-ray pulsars. The subsequent discovery of X-ray eclipses made Swift\\,J1749.4$-$2807 the first eclipsing accreting millisecond X-ray pulsar.} {We obtain additional information on \\src\\ and other burst-only sources. } {We report on the results of a monitoring campaign on the source, carried out for about two weeks with the \\swift,\\ \\inte,\\ and \\rxte\\ satellites.} {The observations showed that the X-ray spectrum (energy range 0.5--40\\,keV) of \\src\\ during the entire event was accurately modeled by an absorbed power-law model ($N_{\\rm H}$$\\simeq$3$\\times$10$^{22}$~cm$^{-2}$, $\\Gamma$$\\simeq$1.7).  X-ray eclipses were also detected in the \\swift\\ data and provides a clear evidence of a dust-scattering halo located along the line of sight to the source. Only one type-I X-ray burst was observed throughout the two-weeks long monitoring. The X-ray flux of \\src\\ decayed below the detection threshold of \\swift\\,/XRT about 11~days after the discovery, in a exponential fashion (e-folding time of $\\tau$=12$^{+7}_{-3}$~days).} {We compare the properties of the outburst observed from \\src\\ with those of the previously known millisecond X-ray pulsars and other transient low mass X-ray binaries.} ",
        "introduction": "\\label{sec:intro}   Low mass X-ray binaries (LMXBs) consist of a low-mass donor star ($<$1~M$_{\\odot}$) and a compact object that accretes matter through an accretion disk. Most LMXBs are transients, i.e. they undergo week-to-month long outbursts during which the accretion takes place at high rates, giving rise to typical (``persistent'') X-ray luminosity of $\\sim$10$^{36-38}$~erg/s. This is $\\gtrsim$100 times higher than the X-ray luminosity displayed in quiescence \\citep[see e.g.] [for reviews]{bildsten98,campana98,liu07}. Depending on the peak X-ray luminosity reached during the outburst ($L_{\\rm peak}$), transient LMXBs are historically classified as ``bright transient'' ($L_{\\rm peak}$$\\simeq$10$^{37-38}$~erg/s), ``faint transient'' ($L_{\\rm peak}$$\\simeq$10$^{36-37}$~erg/s), or ``very faint transients'' \\citep[$L_{\\rm peak}$$\\leq$10$^{36}$~erg/s, see e.g.,][]{wijnands06,campana09}. The outbursts of transient LMXBs are usually interpreted in terms of disk instability models \\citep[see e.g.,][]{frank02}. In a number of these sources, an accreting neutron star (NS) as a compact object has been unambiguously identified by the detection of type-I X-ray bursts and/or coherent pulsations. The former are bright ($\\sim$10$^{38}$~erg/s) and short ($\\sim$10-60~s) flares that result from thermonuclear explosions occurring in the material accreted at the NS surface \\citep[see e.g. ][for reviews]{lewin93,strohmayer06}.  A peculiar subclass of transient LMXBs was discovered about 10 years ago, mainly thanks to the long-term monitoring of the Galactic center carried out with the \\beppo\\ wide field cameras \\citep[WFC,][]{heise99,cocchi01,cornelisse02,cornelisse04}. At odds with the previously known transient LMXBs, these systems appeared to be characterized by no persistent emission before and after the occurrence of a type-I X-ray burst \\citep{cornelisse02}. The upper-limits inferred for their persistent emission were significantly lower than the persistent emission of the bright transients (10$^{35-36}$~erg/s) but not very tight due to the modest sensitivity of the \\beppo\\,-WFCs. For this reason, they were collectively termed ``burst-only sources''. Only in a few cases, a few ks-long \\chan\\ observations carried out after the discovery of a type-I X-ray burst from some of these sources, were able to reveal the presence of a weak persistent emission at a position consistent with that of the X-ray burst. Typical (quiescent) luminosities derived from these detections were $\\sim$10$^{32}$-10$^{33}$~erg/s \\citep{cornelisse02b}.  Owing to their peculiar behavior and the discovery of several new members of this class \\citep[see e.g.,][]{chelovekov07,degenaar09}, the burst-only sources have attracted increasing interest in the past few years \\citep[see also][]{delsanto10}. Systems undergoing type-I X-ray bursts when subject to very low accretion rates allow us to test models of thermonuclear burning in a regime that is still poorly explored \\citep[see e.g.,][]{cooper07,peng07}. Moreover, the low persistent luminosity in quiescence and the spectral properties during the outbursts led to the suggestion that the burst-only sources could be linked to very faint transients and accreting millisecond X-ray pulsars \\citep[AMSPs,][]{wijnands06,campana09,trap09}.  In this paper, we report on the recent discovery of an intense X-ray activity from the burst-only source \\src.\\ We studied the timing and spectral properties of the source during this event by exploiting all the available \\swift,\\ \\inte,\\ and \\rxte\\ data, and report on the detection of millisecond pulsations in the X-ray emission of this source. The results presented here provide a strong support in favor of the association between burst-only sources and AMSPs. In Sect.~\\ref{sec:source}, we describe previous observations of \\src,\\ and in Sect.~\\ref{sec:results} our data analysis and results. Our discussion and conclusions are presented in Sect.~\\ref{sec:discussion}.   \\subsection{SWIFT~J1749.4$-$2807} \\label{sec:source}  \\begin{table*} \\scriptsize \\caption{Observation log of \\src.\\ EXP indicates the exposure time in ks of each observation (for \\inte, it is given separately for ISGRI and JEM-X2 in parentheses), $N_{\\rm H}$ the absorption column density in units of 10$^{22}$~cm$^{-2}$, and $\\Gamma$ the power-law photon index of the X-ray spectrum. For \\inte, an exponential cutoff energy was introduced in the fit (energy fixed at 20~keV). $F_{\\rm obs}^{c}$ is the observed flux and is given in the 0.5-10\\,keV energy band for \\swift\\,/XRT, and in the 3-20~keV energy band for the \\rxte\\,/PCA and \\inte\\,/ISGRI+JEMX-2 in units of $10^{-10}\\,\\mathrm{erg/s/cm^{2}}$.} \\begin{tabular}{@{}llllllll@{}} \\hline \\noalign{\\smallskip} \\swift\\ & & & &\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} OBS ID & START TIME & STOP TIME & EXP & $N_{\\rm H}$ & $\\Gamma$ & $F_{\\rm obs}$ & $\\chi^2_{\\rm red}$/d.o.f. \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 00031686001 &  2010-04-11 23:39:16 & 2010-04-12 04:48:49 &  4.9  & 3.2$\\pm$0.3  &  2.6$\\pm$0.2   & 3.0$\\pm$0.2&  0.97/140 \\\\ \\noalign{\\smallskip} 00031686002 &  2010-04-12 23:54:41 & 2010-04-14 01:32:48 &  2.6  &  3.0$\\pm$0.4  &  2.3$\\pm$0.2 & 2.6$\\pm$0.3  &  0.99/85 \\\\ \\noalign{\\smallskip} 00031686002$^a$  &  2010-04-12 23:54:35 & 2010-04-14 01:31:11 &  0.07  &  4.1$^{+2.6}_{-1.7}$  &  2.4$^{+1.0}_{-0.8}$ &  2.3$^{+0.3}_{-2.3}$  &  (36.6/39)$^c$ \\\\ \\noalign{\\smallskip} 00031686003 &  2010-04-16 16:07:00 & 2010-04-16 18:10:00 &  3.1 & 2.9$\\pm$0.3  &  2.2$\\pm$0.2  & 1.3$\\pm$0.1  &  0.90/87 \\\\ \\noalign{\\smallskip} 00031686004 &  2010-04-17 11:08:00 & 2010-04-17 13:10:00 &  2.9 & 2.7$\\pm$0.5  &  1.9$\\pm$0.2 & 9.0$\\pm$1.0  & 0.99/87  \\\\ \\noalign{\\smallskip} 00031686005 &  2010-04-18 03:34:00 & 2010-04-18 07:07:00 &  3.0 & 2.7$\\pm$0.4  &  1.8$\\pm$0.2  & 5.7$\\pm$0.4  &  0.85/78  \\\\ \\noalign{\\smallskip} 00031686006 &  2010-04-20 16:29:00 & 2010-04-20 23:04:00 &  1.1 & 1.0$^{+3.0}_{-1.0}$  &  1.0$^{+1.1}_{-0.8}$   & $(1.0^{+0.3}_{-1.0})\\times 10^{-1}$  & (15.0/13)$^c$  \\\\ \\noalign{\\smallskip} 00031686007 &  2010-04-21 16:19:00 & 2010-04-21 18:26:00 &  1.6 & 2.7 (fixed)  &  1.8 (fixed)   & $(6.5^{+2.3}_{-2.3})\\times 10^{-3}$ & -    \\\\ \\noalign{\\smallskip} 00031686008$^{b}$ &  2010-04-23 18:25:00 & 2010-04-23 23:38:00 &  3.3 & 2.7 (fixed)  & 1.8 (fixed)   & $<3.3\\times 10^{-3}$ & -  \\\\ \\noalign{\\smallskip} 00031686009$^{b}$ &  2010-04-24 20:23:00 & 2010-04-24 23:21:00 &  2.2 & 2.7 (fixed)  & 1.8 (fixed)  &  $<6.5\\times 10^{-3}$ & -   \\\\ \\noalign{\\smallskip} 00031686010$^{b}$ &  2010-04-25 21:32:00 & 2010-04-25 23:39:00 &  1.0 & 2.7 (fixed)  & 1.8 (fixed)  &  $<8.8\\times 10^{-3}$ & -   \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\rxte\\ & & & &\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 95085-09-01-00 & 2010-04-14 20:59:12 & 2010-04-14 21:34:40 & 1.5 & $3.0 \\pm 1.0$ & $1.8 \\pm 0.1$ & $2.54_{-0.06}^{+0.03} $ & 1.2/43 \\\\ \\noalign{\\smallskip} 95085-09-01-01 & 2010-04-15 16:49:20 & 2010-04-15 17:42:40 & 2.1 & $3.2 \\pm 1.0$ & $1.8 \\pm 0.1$ & $2.21_{-0.05}^{+0.03} $ & 1.3/43 \\\\ \\noalign{\\smallskip} 95085-09-01-02 & 2010-04-15 21:31:12 & 2010-04-16 00:01:04 & 5.7 & $3.6 \\pm 0.6$ & $1.72 \\pm 0.04$ & $2.13_{-0.03}^{+0.02} $ & 1.2/43 \\\\ \\noalign{\\smallskip} 95085-09-02-00 & 2010-04-16 14:46:24 & 2010-04-16 15:41:52 & 2.5 & $3.1 \\pm 1.4$ & $1.7 \\pm 0.1$ & $2.15_{-0.08}^{+0.04} $ & 1.3/43 \\\\ \\noalign{\\smallskip} 95085-09-02-01 & 2010-04-16 16:32:16 & 2010-04-16 17:33:36 & 2.2 & $2.9 \\pm 1.5$ & $1.7 \\pm 0.1$ & $2.06_{-0.09}^{+0.04} $ & 1.3/43 \\\\ \\noalign{\\smallskip} 95085-09-02-02 & 2010-04-16 22:37:20 & 2010-04-17 00:24:32 & 2.6 & $3.4 \\pm 0.4 $ & $1.7 \\pm 0.1$ & $2.07_{-0.09}^{+0.04} $ & 1.0/43 \\\\ \\noalign{\\smallskip} 95085-09-02-09 & 2010-04-17 19:20:16 & 2010-04-17 23:21:36 & 8.0 & $3.6 \\pm 1.6 $ & $1.7 \\pm 0.1$ & $1.23_{-0.06}^{+0.03} $ & 1.4/43 \\\\ \\noalign{\\smallskip} 95085-09-02-03 & 2010-04-18 15:25:20 & 2010-04-18 17:04:32 & 2.7 & 3 (fixed) & $1.6 \\pm 0.1$ & $0.96_{-0.03}^{+0.03} $ & 0.9/44 \\\\ \\noalign{\\smallskip} 95085-09-02-06 & 2010-04-19T07:18:24 & 2010-04-19T07:32:00 & 0.7 & 3 (fixed) & $1.6 \\pm 0.2$ & $0.71_{-0.09}^{+0.06}$ & 1.5/44 \\\\ \\noalign{\\smallskip} 95085-09-02-04 & 2010-04-19 19:31:12 & 2010-04-19 22:24:32 & 6.4 & 3 (fixed) & $1.6 \\pm 0.1$ & $0.51_{-0.04}^{+0.03} $ & 1.3/44 \\\\ \\noalign{\\smallskip} 95085-09-02-07 & 2010-04-20T08:12:16 & 2010-04-20 08:42:56 & 1.0 & 3 (fixed) & $1.5 \\pm 0.2$ & $0.40_{-0.11}^{+0.04} $ & 1.0/44 \\\\ \\noalign{\\smallskip} 95085-09-02-05 & 2010-04-20 16:01:20 & 2010-04-20 19:15:28 & 3.2 & 3 (fixed) & $1.7 \\pm 0.2$ & $0.27_{-0.07}^{+0.02} $ & 1.0/46 \\\\ \\noalign{\\smallskip} 95085-09-02-11 & 2010-04-20 23:10:24 & 2010-04-20 23:28:32 & 0.8 & 3 (fixed) & 1.7 (fixed) & $<1.3\\times 10^{-1}$ & - \\\\ 95085-09-02-08$^d$ & 2010-04-21 13:58:24 & 2010-04-21 14:55:28 & 2.4 & - & - & - & - \\\\ 95085-09-02-10$^d$ & 2010-04-21 15:29:20 & 2010-04-21 17:16:32 & 3.2 & - & - & - & - \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\inte\\ & & & &\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 0720001/0018$^e$ & 2010-04-07 13:31:59 & 2010-04-07 17:13:34 & 15(9) & - & - &$<3$ & \\\\ \\noalign{\\smallskip} 0720001/0019 & 2010-04-10 22:00:00 & 2010-04-11 01:41:36 & 8.4(6.5) & 3 (fixed) & $1.3\\pm0.3$  & $3.3_{-1.3}^{+0.6}$ & 0.5/16 \\\\ \\noalign{\\smallskip} 0720001/0020 & 2010-04-13 15:00:00 & 2010-04-13 18:41:45 & 8.9(11.3) & 3 (fixed) & $1.6\\pm0.4$  & $2.8_{-1.1}^{+0.6}$  & 0.3/15 \\\\ \\noalign{\\smallskip} 0720001/0021$^e$ & 2010-04-21 20:24:43 & 2010-04-22 00:06:28 & 6.9(6.7) & - & - & $<4$ & \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{8}{l}{$^{a}$: This \\swift\\,/XRT observation was carried out in WT mode. $^{b}$: 3$\\sigma$ upper limit. $^{c}$: In these cases this value is C-stat/d.o.f.} \\\\ \\multicolumn{8}{l}{$^{d}$: Observations used for the Galactic diffuse emission estimation.}\\\\ \\multicolumn{8}{l}{$^{e}$: These are $3\\sigma$ upper limits on the source X-ray flux (3--20~keV) estimated from the JEM-X2 mosaics.}\\\\ \\end{tabular} \\label{tab:log} \\end{table*}  \\src\\ was discovered in 2006 during a bright type-I X-ray burst that was initially recorded as a potential gamma ray burst \\citep[GRB060602B,][]{wijnands09}. \\swift\\,/XRT started to follow-up the evolution of the source from 83\\,s after the BAT trigger, and monitored the source outburst for the following 8~d. The source X-ray flux decayed in an approximately power-law fashion (with index $\\sim$-1), fading below the detection threshold of \\swift\\,/XRT in less than $10^6$\\,s. The BAT spectrum extracted at the peak of the type-I X-ray burst could be described by a black-body (BB) model (kT=2.9$^{+0.4}_{-0.3}$~keV) and provided an upper limit to the source distance of $6.7\\pm1.3$\\,kpc \\citep[by assuming that the peak X-ray luminosity of the burst corresponded to the Eddington value, uncertainties at 90\\% c.l.;][]{wijnands09}. During the outburst, the XRT spectrum could be described well by an absorbed power-law model. The measured column density was consistent with being constant throughout the event at a value $N_{\\rm H}$=4$\\times$10$^{22}$~cm$^{-2}$, while the photon index was observed to decrease from $2.7^{+1.5}_{-1.1}$ to $0.5\\pm1.3$ \\citep{campana09}. The estimated X-ray luminosity was $\\sim$10$^{36}$~erg/s at the peak ot the outburst and $\\sim$10$^{32}$~erg/s in the latest \\swift\\ observation available.  \\citet{wijnands09} also reported on three serendipitous detections of \\src\\ in archival \\XMM\\ observations. The first observation was carried out on 2000 September 23, and the second two were performed on 2006 September 22 and 26. The count-rate of the source in the three \\XMM\\ observations was too low to extract any meaningful spectral information. By assuming an absorbed power-law model with $\\Gamma$=2 and $N_{\\rm H}$=3$\\times$10$^{22}$~cm$^{-2}$, \\citet{wijnands09} estimated a 2-10~keV unabsorbed flux of $\\sim$(1-2)$\\times$10$^{-13}$~erg/cm$^{2}$/s. This corresponds to a luminosity of (3-6)$\\times$10$^{32}$~erg/s (assuming a distance to the source of 7~kpc) and is compatible, to within the errors, with that measured about 6~days after the outburst discovered with \\swift.\\  \\src\\ was detected again in a high X-ray luminosity state on 2010 April 10, during the \\textsl{INTEGRAL} Galactic bulge monitoring program \\citep{pavan10,kuulkers07}. \\swift\\ and RXTE target of opportunity observations (ToO) were immediately requested, and these monitored the source outburst for about two weeks. \\rxte\\ observations detected coherent pulsations in the X-ray emission of the source at 518\\,Hz and its second harmonic \\citep{altamirano2010a,bozzo2010}. From these data, a preliminary orbital solution was first obtained by \\citet{belloni10}, and then refined by \\citet{strohmayer2010} by using a pulse phase-coherent technique. The latter authors derived an orbital period of $31740.345\\pm 0.04$\\,s (8.8168 hr), a projected semi-major axis of $a \\sin(i) = 1899.53 \\pm 0.01$\\,lt-ms, and a barycentric frequency for the second harmonic (i.e., twice the spin frequency) of $1035.840025$\\,Hz $\\pm0.4\\,\\mu$Hz. The time of the ascending node was $2455301.1522672 \\pm0.0000014$\\,JD (TDB). This solution implied a mass function of 0.05463$\\pm$0.00018\\,$M_{\\odot}$ and a minimum mass for the companion of 0.475~$M_{\\odot}$ (assuming a NS of 1.4$M_{\\odot}$). The discovery of an X-ray eclipse in the \\rxte\\ light curve was reported by \\citet{markwardt10a}; this is the first eclipse observed from an accreting millisecond X-ray pulsar. The most accurate source position to date was provided by \\citet{yang10} at $\\alpha_{\\rm J2000}$=17$^{\\rm h}$49$^{\\rm m}$31$\\fs$80 and $\\delta_{\\rm J2000}$=-28${\\degr}$08$\\arcmin$04$\\farcs$9, with a 90\\% confinement radius of 1.9$\\farcs$, based on \\swift\\,/XRT observations. ",
        "conclusions": "\\label{sec:discussion}  We have reported on the monitoring of the burst-only source \\src,\\ which was discovered to undergo a new X-ray outburst by \\inte\\ on 2010 April 10.  \\subsection{The outburst decay} The peak outburst luminosity (see Fig.~\\ref{fig:lcurve}) was $\\sim$1.8$\\times$10$^{36}$~erg/s (assuming a distance of 7~kpc), only slightly higher than that reported for the outburst that occurred in 2006 ($\\sim$10$^{36}$~erg/s). Despite this similarity in peak X-ray luminosity, the evolution in the source X-ray flux during the two events was significantly different. The 2006 outburst displayed a clear power-law-like decay \\citep[index $\\sim$-1,][]{wijnands09}, and lasted less than $\\sim$6~days (see also Sect.~\\ref{sec:source}). In contrast, the 2010 outburst lasted more that 11~days\\footnote{Note that we were unable to determine the exact starting point of the outburst, since the first available observation prior to the discovery of the renewed activity is on 2010 April 7 and the relatively short exposure time enables to derive only a poorly constraining upper limit (see Table~\\ref{tab:log}).} and was characterized by an exponential decay (see Sect.~\\ref{sec:results}).  Disk instability models for the outbursts of transient NS LMXBs predict that both linearly and exponentially decaying outburst can be produced \\citep{king98}. \\citet{king98b} showed that, when the irradiation of the accretion disk from the X-ray emission of the NS is strong enough (i.e., the peak X-ray luminosity at the onset of the outburst is high), the disk is completely ionized out to its outer edge and the light curve of the outburst follows an exponential decay. In this phase, most of the disc mass can be accreted onto the NS in a viscous timescale. A change in the profile of the outburst is expected once the source has faded below a certain luminosity level. Below this threshold, only part of the accretion disk can be ionized. As a consequence, both the mass accretion rate onto the NS and the outer boundary of the ionized region in the disk begin to decrease with time, leading to a linear (rather than exponential) decay \\citep{king98,shahbaz98,powell07}.  The value of the critical X-ray luminosity at which these changes occur depends on the properties of the system (total mass and orbital period). In the case of \\src\\,, this luminosity is expected to be on the order of 10$^{36}$~erg/s  \\citep{shahbaz98}, thus comparable to the peak X-ray luminosity estimated during both the events in 2006 and 2010. As the event in 2010 was characterized by an X-ray luminosity slightly higher than that measured at the peak of the outburst in 2006, it is possible that this relatively small difference was sufficient to cause a different behavior during the outburst decay. We note that, even though the peak luminosity of the two outbursts was almost comparable, during the outburst that occurred in 2010 the source remained in a high X-ray luminosity state for a significantly longer time interval than in 2006. This might have caused the ionization of a larger portion of the accretion disk and thus led to an exponential outburst decay. Given the relatively poor number of \\swift\\ and \\rxte\\ observations during the end of the outburst in 2010, it is not possible to investigate more deeply whether if a switch between exponential and linear decay occurred at some point during the steep decay of the 2010 outburst.  The discovery of millisecond pulsations in the X-ray emission of \\src\\ secured the previously suggested association of this source (and other burst-only sources in general) with the class of the accreting AMSPs. This permits a comparison between the two classes. In particular, AMSPs are known to undergo week-to-month-long outbursts that usually exhibit an initial exponential decay, followed by a linear one, in agreement with the prediction of the disk instability model discussed above. An in-depth study of the outburst decay of a number of AMSPs was carried out by \\citet{powell07}. These authors showed that most of these sources display a ``knee'' in their decay, when the X-ray luminosity falls below the threshold at which the switch between exponential to linear decay is predicted.  This behavior was observed in the AMSPs SAX\\,J1808.4$-$3658, XTE\\,J1751$-$305, and XTE\\,J0929$-$314. Some anomalies in the outburst decay were reported for the AMSPs SAX\\,J1808.4$-$3658, XTE\\,J1807$-$294, and IGR\\,J1751$-$305. During an outburst in 2000, SAX\\,J1808.4$-$3658 displayed several unexpected ``re-flares'' about $\\sim$20~days after the initial exponential+linear decay \\citep{wijnands06,campana08}, whereas the source XTE\\,J1807$-$294 underwent in 2003 a rather long ($\\gtrsim$100~day) pure exponential outburst without any clear evidence of a linear phase \\citep[see e.g.,][and references therein]{falanga2005,powell07}. Another AMSP, IGR\\,J1751-305, displayed a few outbursts lasting exceptionally only a few days and reaching a relatively low X-ray luminosity \\citep[a few mCrab, see e.g.,][]{falanga07,markwardt07}. The outbursts observed from \\src\\ in 2006 and 2010 were both significantly shorter than those observed from the other AMSPs, and only the anomalous outbursts detected from IGR\\,J1751-305 resemble those of \\src\\ .  Similar dim outbursts have been observed from the very faint X-ray transients (VFXTs, see Sect.~\\ref{sec:intro}), and an association between these sources and the burst-only sources was suggested by \\citet{campana09}. Different ways of interpreting the behavior of the VFXTs were reported by \\citet{king06}. These authors showed that the VFXTs hosting NS primaries can be relatively well understood if the material accreted on these systems originates from brown dwarfs or planetary companions. However, in the case of \\src\\ such low-mass companions is ruled out by the mass function of the system, which implies a minimum companion mass of $\\sim$0.5$M_{\\odot}$ assuming a NS of  $1\\,M_{\\odot}$.  \\subsection{Timing features} In the case of \\src\\,, an in-depth study of the properties of the primary and secondary stars is possible for the first time by using the orbital solution and characteristics of the X-ray eclipses. At present, \\src\\ is the only eclipsing AMSPs. An extensive discussion of these issues is beyond the scope of the present work, and will be reported in a separate paper (Campana et al., 2010, in preparation).  In Sect.~\\ref{sec:halo}, we carried out a careful analysis of the residual X-ray emission of \\src\\ during the eclipse. While our analysis of the \\rxte\\ data could not unambiguously identify the origin of this emission, we found in the \\swift\\ data strong evidence that residual emission during the eclipse is caused by a scattering halo in the direction of the source. The effect of a dust scattering halo is most prominent when a source is in eclipse and at lower energies ($\\lesssim$3~keV, see Sect.~\\ref{sec:halo}). Since \\src\\ is the first eclipsing AMSP, it is also the first source in this class for which this effect could be detected.  An additional peculiarity of \\src\\ was revealed by the Fourier analysis of the \\rxte\\ data, which showed a signal at the spin frequency of the NS and its second harmonic. This is consistent with the fact that the pulse profiles clearly displayed a double peak (see Fig.~\\ref{fig:timepulses}). This feature has never been observed in another AMSP and is most likely due to emission from both accreting polar caps that sweeps across our line of sight. This geometry is possible because the system is seen nearly edge-on (as confirmed by the detection of X-ray eclipses). Moreover, the study of the pulse profile in different \\rxte\\ observations has found that the shape and amplitude of the two peaks changed significantly during the outburst decay. We propose that this behavior is caused by changes in the geometry of the emitting region close to the NS surface  and/or changes in the radius and/or thickness of the inner edge of the accretion disk.  \\subsection{Type I X-ray burst} Our analysis of the data presented here has also detected a type-I X-ray burst in the JEM-X2 light curve. This is the second type-I X-ray burst reported so far from \\src.\\ The first burst was discovered with \\swift\\,/BAT \\citep{wijnands09} and was characterized by a peak flux a factor of $\\sim$2 higher than that measured by JEM-X2 (see also Sect.~\\ref{sec:source}). The statistics of the JEM-X2 were of too low quality to perform a time-resolved spectral analysis of the burst, and no evidence of a photospheric radius expansion could be found (see Sect.~\\ref{sec:results}). The observation of a type-I X-ray burst can be used to derive an upper limit to the source distance by assuming that the peak X-ray luminosity of the burst corresponded to the Eddington value $L_{\\rm Edd}$$\\approx$3.8$\\times$10$^{38}$~erg/s \\citep[as empirically derived by][for a helium burst]{kul03}. However, since the peak flux of the JEM-X2 burst was significantly lower than that observed previously by \\swift,\\ the latter already provided the most restrictive upper limit\\footnote{We note however that, if a pure hydrogen burst is considered, then the inferred upper limit to the source distance would be $\\approx$5~kpc.} of 7~kpc on the distance to \\src\\ (see Sect.~\\ref{sec:intro}). Using this distance, the persistent unabsorbed 0.1-100~keV flux of the source at the time of the JEM-X2 burst would translate into a bolometric luminosity of $L_{\\rm pers}\\approx~4.7\\times 10^{36}$~erg/s, or 1.2\\% $L_{\\rm Edd}$. This corresponds to a local accretion rate per unit area of $\\dot{m}\\approx$3.2$\\times$10$^{3}$~g/s/cm$^2\\simeq$1.2\\%$\\dot m_{\\rm Edd}$, where we have used the relation $L_{\\rm pers}=4\\pi R^2\\dot m(GM/R)/(1+z)$ \\citep[with $z$=0.31 the NS gravitational redshift; see e.g.,][]{lewin93}.  Theoretical models predict that, when 0.01$<$$\\dot{m}$/$\\dot{m_{\\rm edd}}$$<$0.1, H burns stably through the hot CNO cycle and a layer composed of pure He develops underneath the NS surface. This layer can then ignite by means of the 3$\\alpha$ process and lead to a pure He burst, with a typical rise time of $\\sim$1~s and a duration of $\\sim$10~s. %The rise time of $\\sim1$~s, the total duration of $\\sim10$~s, and the $\\dot{m}$ These quantities, estimated for the burst detected by JEM-X2 (see Sect.~\\ref{sec:results}), are fully compatible with those expected for a pure He burst.  We note that pure He runaways were also reported for the other burst-only source GRS\\,1741.9-2853 \\citep{trap09}. The detection of this kind of type-I X-ray bursts from the burst-only sources is particularly intriguing because it would argue against the idea that these sources are the prototypes of the (poorly observed) H-burning bursts with low accretion rates \\citep[see Sect.~\\ref{sec:intro} and, e.g.][]{peng07}.  Finally, we can check the consistency of the results derived above for the type-I X-ray burst by evaluating the theoretically expected recurrence time of the burst. We first estimate the ignition depth of the burst, $y_{\\rm ign}$, using the equation $E_{\\rm burst}=4\\pi R^2y_{\\rm ign}Q_{\\rm nuc}/(1+z)$, where $E_{\\rm burst}=4\\pi d^2f_b=2.0\\times10^{39}$ erg ($d$/7 kpc), $f_b$ is the measured fluence of the burst (see Table~\\ref{tab:burst}), and  $Q_{\\rm nuc}\\approx 1.6$ MeV corresponds to the nuclear energy release per nucleon for complete burning of helium to iron group elements \\citep{wallace81,fujimoto87}. We obtained $y_{\\rm ign}$=$1.4\\times$10$^8$~${\\rm g\\ cm^{-2}}$. For the above values of the local accretion rates and ignition depth, the expected recurrence time of a He bursts is about $\\Delta$t=($y_{\\rm ign}$/$\\dot m$)$(1+z)$$\\simeq$0.6~days (independent of the assumed distance). The burst detected by JEM-X2 occurred $\\sim6.8$~d after the first available \\inte\\ observation (see Fig.~\\ref{fig:lcurve}). However, the total effective exposure time on the source was $\\sim$0.9~days, thus compatible with the observation of a single type-I X-ray burst throughout the observational period.\\\\   At the time of writing this work, two other papers were submitted describing the same source: \\citet{markwardt10b} and \\citet{altamirano2010b}. These authors focused mostly on the timing analysis of the system and discussed in detail the dynamical constraints on the binary system that could be derived by using the eclipses detected in the RXTE data. In our paper, we have concentrated more on the study and interpretation of the evolution of the timing and spectral properties of the source during the outburst. We checked (where possible) that the results derived in our paper are fully compatible with those reported in the other two papers. We have also discussed for the first time the properties of the eclipse found in the \\swift\\ data, and provided an interpretation of the residual flux during the eclipse."
    },
    "1005/1005.0878_arXiv.txt": {
        "abstract": "We compare non-locality of interactions between different scales in hydrodynamic (HD) turbulence and magnetohydrodynamic (MHD) turbulence in a strongly magnetized medium. We use 3-dimensional incompressible direct numerical simulations to evaluate non-locality of interactions. Our results show that non-locality in MHD turbulence is much more pronounced than that in HD turbulence. Roughly speaking, non-local interactions count for more than 10\\% of total interactions in our MHD simulation on a grid of $512^3$ points. However, there is no evidence that non-local interactions are important in our HD simulation with the same numerical resolution. We briefly discuss how non-locality affects energy spectrum. ",
        "introduction": "Turbulence is commonly observed in astrophysical fluids and in many cases such turbulence is accompanied by a strong magnetic field, which has a large impact on the dynamics of the turbulent cascade. Since turbulence influences many astrophysical processes (e.g.~transport of mass and angular momentum, star formation, fragmentation of molecular clouds, heat and cosmic ray transport, magnetic reconnection, etc.), understanding scaling properties of magnetohydrodynamic (MHD) turbulence is essential for theoretical astrophysics. For this reason, rich literature exists regarding scaling relations of MHD turbulence (See Goldreich \\& Sridhar 1995; Biskamp 2003 and references therein; see also Cho \\& Vishniac 2000b; Maron \\& Goldreich 2001; M\\\"uller, Biskamp, \\& Grappin 2003; M\\\"uller \\& Grappin 2005; Boldyrev 2005, 2006; Beresnyak \\& Lazarian 2006; Mason, Cattaneo \\& Boldyrev 2006; Gogoberidze 2007; Matthaeus et al. 2008).  In hydrodynamic (HD) turbulence, energy cascades down to smaller scales. Kinetic energy contained in an ``eddy'' is transferred to smaller eddies by shearing motions of other eddies (see, for example, Frisch 1995). Most theories on turbulence assume locality of interactions, which means interactions between similar size eddies dominate in such energy cascade. In Fourier space, this means that a Fourier mode at a wavenumber $k=|{\\bf k}|$, where ${\\bf k}$ is the wavevector, interacts mainly with other modes having similar wavenumbers and transfers its energy to modes that have larger wavenumbers. Recently many researchers have investigated locality in HD turbulence (Mininni, Alexakis, \\& Pouquet 2008; Alexakis, Mininni, \\& Pouquet 2007; see also Verma et al. 2005).  In MHD turbulence with a strong mean field (${\\bf B_0}$), locality is also generally assumed. However, in the MHD case, the nature of energy cascade is slightly different. % In the incompressible limit, any magnetic perturbation propagates {\\it along} the magnetic field line. To the first order, the speed of propagation is constant and equal to the Alfv\\'en speed $V_A=B_0/\\sqrt{4\\pi \\rho}$, where $\\rho$ is the density. Since wave packets are moving along the magnetic field line, there are two possible directions for propagation. If all the wave packets are moving in one direction, then they are stable to nonlinear order (Parker 1979). Therefore, in order to initiate turbulence, there must be opposite-traveling wave packets and the energy cascade occurs only when they collide. Therefore, in the MHD case, locality means that a wave packet (or ``eddy'') transfers energy to smaller scale wave packets by shearing motions of opposite-traveling wave packets of similar size.  There have been some discussions about non-locality in MHD turbulence with a strong mean field.\\footnote{ When the mean field is weak or zero, turbulence structure is very different (see for example Cho et al. 2009). There are many discussions about non-locality in this regime % (see for example Alexakis, Mininni, \\& Pouquet 2005b; Lessinnes, Verma, \\& Carati 2008; Yousef et al. 2009; Aluie \\& Eyink 2010). } For example, Alexakis (2007) theoretically studied non-local model of MHD turbulence. In their inspiring work, Beresnyak \\& Lazarian (2010) numerically studied non-locality and argued that ``MHD turbulence is fairly non-local, at least less local than hydrodynamic turbulence'' (see also Beresnyak \\& Lazarian 2009). They claimed that `` a) the lack of visible bottleneck effect in MHD turbulence, while it is clearly present in hydro turbulence, and b) the dependence of kinetic and magnetic spectra on driving'' support this idea. Teaca et al. (2009) calculated anisotropic energy transfer in Fourier space. But they did not pay much attention to the locality issue.   In this paper, we quantitatively evaluate non-locality of HD and MHD turbulence and present a direct evidence that non-locality is clearly present in MHD turbulence. We consider only balanced strong MHD turbulence. Here balanced MHD turbulence means that amplitudes of two opposite-traveling wave packets are almost equal. In \\S2, we describe our numerical setup. In \\S3, we present our results. In \\S4, we briefly discuss how non-locality affects energy spectrum  and give summary. ",
        "conclusions": "When shearing motions of the outer scale eddies influence energy transfer of inertial range eddies, energy spectrum becomes flatter than the Kolmogorov one. Suppose that the shearing motions of the outer scale eddies completely dominate energy cascade. In this case, from \\be Z_l^2/t_{cas} \\sim Z_l^2/(L/v_L) \\propto Z_l^2 = constant, \\ee we can easily show that energy spectrum is $E(k)\\propto k^{-1}$, where $Z_l$ is an Els\\\"asser variable at scale $l$, $L$ the outer scale, and $v_L$ the rms velocity at the outer scale (see Equation (1) of Cho, Lazarian, \\& Vishniac 2003). If the shearing motions of the outer scale eddies do not completely dominate, we will have a spectrum between $k^{-1}$ and $k^{-5/3}$.   Indeed, in Figure~\\ref{fig_t} we observe that $E_v(k)+E_b(k)$ (hence spectrum of ${\\bf Z}^+$ or ${\\bf Z}^-$) in MHD is flatter than the Kolmogorov spectrum. This is consistent with earlier numerical results (see, for example, Maron \\& Goldreich 2001; M\\\"uller et al. 2003).    In summary, we have found the following results. \\begin{enumerate} \\item We have developed a quantitative method to measure non-locality (see Figure~\\ref{fig_ratio}(b) or Figure~\\ref{fig_pq}).  \\item Our numerical calculations show non-locality is more pronounced in MHD turbulence than in HD turbulence. This result confirms an earlier finding by Beresnyak \\& Lazarian (2010).  \\item There are two forms of non-locality in MHD turbulence: non-local influence of shearing motions, which does {\\it not} involve energy transfer between different scales, and non-local energy transfer between different scales (\\S\\ref{sect:et}). In MHD, the former type of non-locality (i.e. non-local influence of outer scale shearing motions) is more important. It is not clear whether the latter type is important.  \\item In MHD, non-locality is not negligible so that it might affect dynamics of turbulent cascade.  \\end{enumerate}"
    },
    "1005/1005.1942_arXiv.txt": {
        "abstract": "{ We present {\\it Herschel}/PACS spectroscopic maps of the dwarf galaxy NC4214 observed in 6 far infrared fine-structure lines: [C~{\\sc ii}] 158$\\mu$m, [O~{\\sc iii}] 88$\\mu$m, [O~{\\sc i}] 63$\\mu$m, [O~{\\sc i}] 146$\\mu$m, [N~{\\sc ii}] 122$\\mu$m, and [N~{\\sc ii}] 205$\\mu$m. The maps are sampled to the full telescope spatial resolution and reveal unprecedented detail on $\\sim$ {{\\revised 150}} pc size scales. We detect [C~{\\sc ii}] emission over the whole mapped area, [O~{\\sc iii}] being the most luminous FIR line. The ratio of [O~{\\sc iii}]/[C~{\\sc ii}] {{\\revised peaks at about 2 toward the sites of massive star formation}}, higher than ratios seen in dusty starburst galaxies. The [C~{\\sc ii}]/CO ratios are 20 000 to 70 000 toward the 2 massive clusters, which are at least an order of magnitude larger than spiral or dusty starbursts, {{\\revised and cannot be reconciled with single-slab PDR models}}. Toward the 2 massive star-forming regions, we find that L$_{\\rm [CII]}$ is 0.5 to 0.8\\% of the L$_{\\rm TIR}$. All of the lines together contribute up to 2\\% of L$_{\\rm TIR}$. These extreme findings are a consequence of the lower metallicity and young, massive-star formation commonly found in dwarf galaxies. {{\\revised These conditions promote}} large-scale photodissociation into the molecular reservoir, {{\\revised which is evident in the FIR line ratios}}. This illustrates the necessity {{\\revised to move to multiphase models}} applicable to star-forming clusters or galaxies as a whole. } ",
        "introduction": "As building blocks of larger galaxies, dwarf galaxies in the early universe presumably played an important role in the evolution of galaxies we see today. Understanding how stars form and evolve under low metallicity conditions and the process of subsequent enrichment of the interstellar medium (ISM) will help us constrain possible scenarios of galaxy formation and evolution. The Local Group dwarf galaxies provide laboratories to study in detail the effects of the lower metal abundance on the process of star formation and the feedback on the ISM.  {\\revised To examine the interplay between massive star formation and low metallicity ISM}, we have mapped the far infrared (FIR) fine structure lines in the Local Group Magellanic-type irregular galaxy, \\object{NGC4214}, as part of the {\\it Herschel} key proposal, SHINING (P.I. E. Sturm). NGC4214 harbors 2 main star-forming complexes containing hundreds of O stars as well as a super star cluster (SSC). Our aim here is to characterise the physical conditions and the structure of the different phases of the ISM in NGC4214. The proximity of NGC4214 \\citep[2.9Mpc,][]{maiz-apellaniz-2002} and its low-metallicity \\citep[$\\rm{log(O/H)+12=8.2}$,][]{kobulnicky-skillman-1996} make it a choice target. The {PACS} spectrometer onboard the {\\it Herschel Space Observatory} allows us to zoom into the low-metallicity photodissociation regions (PDRs) and their surroundings, tracing physical properties from {\\revised 125 to 165} pc size scales.  The tools for this study are the FIR fine-structure lines arising from some of the most abundant species C, O, and N, which provide the most important cooling channels in the neutral and diffuse ionized ISM of galaxies. These lines trace H~{\\sc ii} ~regions and PDRs, which are neutral atomic and molecular regions where penetrating FUV photons (6 eV $< h\\nu <$ 13.6 eV) dominate the energy balance and the chemical composition of the gas. Most of the ISM in galaxies is contained in PDRs. The FIR lines in combination with PDR models \\citep{tielens-hollenbach-1985,wolfire-1990,ferland-1998,kaufman-1999,abel-2005,roellig-2006} can unveil the structure of these regions, which depends on both the strength of the FUV radiation field density ($\\rm{G_{0}}$) illuminating the cloud and the hydrogen density ($\\rm{n_{H}}$). ",
        "conclusions": "\\label{sect:discussion} These observations of NGC4214 with {\\it Herschel}/{PACS} have enabled us to spatially resolve the FIR fine-structure lines in the low metallicity dwarf galaxy NGC4214. A very striking result from the mapping is the brightness of the [C~{\\sc ii}]~158$\\mu$m , [O~{\\sc i}]~63$\\mu$m, and the [O~{\\sc iii}]~88$\\mu$m lines. The [C~{\\sc ii}] alone corresponds to 0.5\\% of the total infrared luminosity. The combined luminosity of the FIR lines reaches 2\\% of the L$_{\\rm TIR}$. The [O~{\\sc iii}] 88$\\mu$m line is more confined to star-forming regions and the [O~{\\sc iii}]/[C~{\\sc ii}] reaches almost a factor of 2 close to the SSC in NGC4214-I.  The {\\revised [C~{\\sc ii}]/CO ratios are also exceptionally high} (20\\,000 $-$ 70\\,000). These extreme ratios are indicative of an ISM that is strongly affected by photodissociation and reflect the combined effects of low metallicity ISM and intense star formation. It should be borne in mind that, when comparing with models, the models predict line ratios for a single, homogeneous PDR. In reality, we observe a multiphase medium that {\\revised probably harbors many extremely} different PDR conditions. Therefore, this preliminary comparison indicates the need for more complete (complex) models to be able to derive accurate physical quantities."
    },
    "1005/1005.1203_arXiv.txt": {
        "abstract": " ",
        "introduction": "The primordial curvature inhomogeneities, generated by the inflationary mechanism, obey nearly Gaussian statistics. The deviations from Gaussianity, while expected to be small, provide a unique window into the physics of inflation. For example, single-field slow-roll models of inflation lead to a small level of non-Gaussianity (NG), so that an observation of a large NG would indicate a deviation from this paradigm.  Until a few years ago, the main tool to constrain NG was considered to be the statistics of the cosmic microwave background (CMB) temperature field, since inhomogeneities at the CMB epoch are small and their physics can be described by a perturbative treatment. In recent years, however, thanks to observations and developments in the theory, the large-scale structure (LSS) of the universe has emerged as a complementary probe to constrain primordial NG. While it is true that the $n$-point functions of the density field on small scales are dominated by the recent gravitational evolution, and do not reflect anymore the statistics of primordial perturbations, it turns out that the abundance of very massive objects, which form out of high peaks of the density perturbations, is a powerful probe of primordial NG. In this context, much attention has been given recently to three possible methods of constraining the magnitude and shape of the primordial NG with the LSS: the galaxy power spectrum, the galaxy bispectrum and the mass function. It was pointed out in \\Cites{Dalal:2007cu, Matarrese:2008nc} that a NG of a local type induces a scale dependence on the galaxy power spectrum, thus making it a sensitive probe of the magnitude of local NG $f_{NL}^{loc}$. From \\Cite{Slosar:2008hx} one finds the following constraints: $-29 < f^{loc}_{NL} < +69$, already comparable with those obtained from CMB measurements in \\Cite{Komatsu:2010fb}: $-10 < f^{loc}_{NL} < +74$. The future is even more promising, with precisions of $\\Delta f^{loc}_{NL} \\sim 10$~\\cite{Sartoris:2010cr} and $\\Delta f^{loc}_{NL} \\sim 1$~\\cite{Carbone:2008iz,Carbone:2010sb,Cunha:2010zz} being claimed for future surveys. The galaxy bispectrum is also a promising probe of NG as it could be more sensitive to other triangle configurations~\\cite{Sefusatti:2009qh}. The mass function -- which is the focus of this work, and which we discuss in detail below -- has been used for example in~\\Cite{Sartoris:2010cr} together with the scale dependent bias to produce forecasts for future surveys, and in~\\Cite{Jimenez:2009us} in an attempt to explain the presence of a very massive cluster at a large redshift as an indication of a large NG. For more references and information we refer the reader to reviews summarizing recent results on these topics~\\cite{Verde:2010wp, Desjacques:2010jw}.  The formation of bound dark matter halos from initially small density perturbations, as seen in numerical simulations, is a complicated and violent process. Some insight into the physics involved has been gained from the study of analytical models. The quantity of interest is the halo mass function, defined as the number density of dark matter halos with a mass between $M$ and $M + dM$, \\be \\frac{dn}{dM} = \\frac{\\bar\\rho}{M^2}f(\\sig)\\left| \\frac{d\\ln\\sig}{d\\ln M} \\right|\\,, \\la{mass-func-conventional} \\ee where $\\bar{\\rho}$ is the average density of the universe, $\\sigma(M)$ is the variance of the density contrast $\\del_R$ filtered on some comoving scale $R$ corresponding to the mass $M$, and the function $f(\\sigma)$ is to be computed. Throughout this work, we will refer to $f(\\sigma)$ itself as the mass function. A very useful tool in the analysis is the spherical collapse model~\\cite{Gunn:1972sv}, which predicts that the value of the linearly extrapolated density contrast of a spherical halo, at the time when the halo collapses, is $\\delc\\simeq1.686$, with a weak cosmology dependence. This value serves as a collapse threshold for determining which inhomogeneous regions will end up as collapsed objects. Using this idea, Press \\& Schechter~\\cite{Press:1973iz} (PS) first computed the mass function $f(\\sigma)$ in the case of Gaussian initial conditions. Their calculation however suffered from a problem of undercounting which affects the overall normalization -- their approach does not count underdense regions embedded in larger overdense regions as eventually collapsed objects. To account for this discrepancy, PS introduced an ad-hoc factor of $2$ by demanding that the mass function be correctly normalized, such that all the mass in the universe must be contained in collapsed objects. In the excursion set approach, Bond \\etal~\\cite{Bond:1990iw} resolved this issue and derived a correctly normalized mass function, for Gaussian initial conditions. They argued that the filtered density contrast $\\del_R$ follows a random walk as a function of the filtering scale, and the problem of computing $f(\\sigma)$ is translated into the problem of finding the rate of ``first crossing'' of the barrier $\\delc$, whose solution is well-known. We will study this formalism in detail in section~\\ref{calc-f} for the more general non-Gaussian case.  Turning to non-Gaussianities, the most popular non-Gaussian mass functions are those due to Matarrese, Verde and Jimenez~\\cite{Matarrese:2000iz} (MVJ) and LoVerde~\\etal~\\cite{LoVerde:2007ri} (LMSV). Both groups used the PS approach, by modifying the probability density function for the (linearized) density contrast to describe non-Gaussian initial conditions. In their prescription, the relevant object is the \\emph{ratio} $R_{\\rm ng}$ of non-Gaussian to Gaussian mass functions. The full mass function is usually taken as the product of $R_{\\rm ng}$ and an appropriate Gaussian mass function as given by $N$-body simulations, e.g. the Sheth \\& Tormen mass function~\\cite{Sheth:2001dp}. It is not clear however that this is the correct way to proceed. Indeed, in a series of papers~\\cite{Maggiore:2009rv,Maggiore:2009rw,Maggiore:2009rx}, Maggiore \\& Riotto (MR) presented a rigorous approach to the first-passage problem in terms of path integrals, and in~\\Cite{Maggiore:2009rx} they pointed out that a PS-like prescription in fact misses some important non-Gaussian effects stemming from $3$-point correlations between \\emph{different} scales (so-called ``unequal time'' correlators).  On the other hand, MR treated non-Gaussian contributions to $f(\\sig)$ by simply linearizing in the $3$-point function of $\\del_R$, i.e. by linearizing in the non-Gaussian parameter \\fnl. Since the NG are assumed to be small, in the sense that the parameter $\\ep = \\langle\\delta^3\\rangle/\\sigma^{3}$ satisfies $\\ep\\ll1$, one might expect that such a perturbative treatment is valid. However, another  crucial ingredient in the problem is that the length scales of interest are large, which leads to a \\emph{second} small parameter $\\nu^{-1}$ where $\\nu = \\delc/\\sig$. This is evident in the calculations of MR, who crucially use $\\nu^{-2}\\propto\\sig^2$ as a small parameter. Any perturbative treatment now depends not only on the smallness of $\\ep$ and $\\nu^{-1}$ individually, but also on the specific combinations of these parameters which appear in the calculation. It is known (and we will explicitly see below) that a natural combination that appears is $\\ep\\nu^3$, which can become of order unity on scales of interest. The mass functions given by LMSV and MR therefore break down as valid series expansions when this occurs. Interestingly, MVJ's PS-like treatment on the other hand involved a saddle point approximation,  allowing them to \\emph{non-perturbatively} account for the $\\ep\\nu^3$ term (which appears in an exponential in their approach). For a discussion, see \\Cite{Valageas:2009vn}.   It appears to us therefore, that there is considerable room for improvement in the theoretical calculation of the mass function. The goal of our paper is twofold. Firstly, we present a rigorous calculation of the mass function in the following way : (a) we use the techniques developed by MR in~\\Cites{Maggiore:2009rv,Maggiore:2009rw,Maggiore:2009rx}, which allow us to track the complex multi-scale correlations involved in the calculation, and (b) we demonstrate that MR's approach can be combined with saddle point techniques (used by MVJ), to non-perturbatively handle terms which can become of order unity.  This leads to an expression for the mass function which is valid on much larger scales than those presented by MR and LMSV.  Secondly, by keeping track of the terms ignored, we calculate theoretical error bars on the expressions for $f(\\sig)$ resulting not only from our own calculations, but also for those of the other authors~\\cite{Matarrese:2000iz,LoVerde:2007ri,Maggiore:2009rx}. Since the terms ignored depend on $\\nu$ in general, these error bars are clearly scale dependent. This allows us to estimate the validity of each of the expressions for the mass function at different scales, but importantly it also allows us to analytically compare between different expressions.  In this paper we will not explicitly account for effects of the ellipsoidal collapse model \\cite{Bardeen:1985tr,Sheth:1999su}, since these are expected to be negligible on the very large scales which are of interest to us. For a recent treatment of ellipsoidal collapse effects in the presence of non-Gaussianities on scales where $\\ep\\nu^3\\ll1$, see Lam \\& Sheth \\cite{Lam:2009nb}. For a different approach to computing the non-Gaussian halo mass function, see \\Cite{Dalal:2007cu}, where the authors proposed that this mass function can be approximated as a convolution of the \\emph{Gaussian} mass function with a probability distribution function that maps between halos identified in Gaussian and non-Gaussian $N$-body simulations. This probability distribution itself was approximated as a Gaussian, with mean and variance fit from simulations. While in practice this approach is easy to implement, it relies heavily on the output of $N$-body simulations. Our approach, on the other hand, allows us to compute a mass function almost entirely from first principles.  This paper is organized as follows. In section~\\ref{models} we fix some notation and briefly introduce the two most popular shapes of primordial NG, i.e. the local and equilateral ones. In section~\\ref{calc-f} we present our calculation of the mass function. In section~\\ref{consistency} we discuss certain subtleties regarding the truncation of the perturbative series, and also compare with the other expressions for $f(\\sig)$ mentioned above. In section~\\ref{diffuse-filter} we discuss the effects induced by some additional complications introduced in the problem due to the specific choice of the filter function~\\cite{Maggiore:2009rv}, which we take to be a top-hat in real space, and due to the inclusion of stochasticity in the value of the collapse threshold $\\delc$  (which is also expected to partially account for effects of ellipsoidal collapse)~\\cite{Maggiore:2009rw} . In section~\\ref{results} we compare our final result~\\eqn{f-fullcorrected} with those of other authors, including  theoretical errors for each, and conclude with a brief discussion of the results and directions for future work. Some technical asides have been relegated to the Appendices. ",
        "conclusions": "\\la{results} \\noindent In this section we conclude with our final results for the non-Gaussian halo mass function, comparing our approach with previous work. In principle, we should compare the full expressions for the mass functions of various authors with ours. However, recall that for MVJ and LMSV one has to multiply an analytically predicted ratio $R_{\\rm ng} = f(\\nu,M,\\fnl)/f(\\nu,M,\\fnl=0)$ with a suitable \\emph{Gaussian} mass function based on fits to simulations. It is not clear how to compute theoretical error bars on the latter. On the other hand, the object $R_{\\rm ng}$ itself is an unambiguous theoretical prediction of every approach, that is MVJ, LMSV, MR and our work, and we can compute theoretical errors on it. In this work, we will restrict ourselves to comparing the different expressions for $R_{\\rm ng}$. In future work, we hope to compare both $R_{\\rm ng}$ and the full mass function with the results of $N$-body simulations.  In~\\fig{final-comparison-z1} and~\\fig{final-comparison-z1k} we plot the ratio $R_{\\rm ng}$, respectively without and with the filter effects, at redshift $z=1$. In this way we can explicitly disentangle the errors due to an approximate treatment of non-Gaussian effects, from those due to the filter effects. We compare our expression \\eqref{f-fullcorrected} with the expressions of MR~\\eqref{f-MR-detailed}, LMSV~\\eqref{f-LMSVlin} and \\eqref{f-LMSVquad}, and of MVJ~\\eqref{f-MVJ}. Notice that, when considering the filter effects, the Gaussian function that enters in the ratio $R_{\\rm ng}$ is defined to be the function with $\\fnl=0$ (i.e. without NG but with filter effects when present). We use the local model, setting $\\fnl=100$ and $\\gnl=0$, and use the reference \\Lam CDM cosmology described in section \\ref{models}. We do not explicitly show the final results for the equilateral model, but they are qualitatively similar. As is commonly done in the literature, we modify the LMSV and MVJ curves by applying the Sheth \\etal\\ correction of $\\del_{c0}\\to\\sqrt{a}\\del_{c0}$. An identical correction is already present in the expressions \\eqref{f-fullcorrected} and \\eqref{f-MR-detailed} due to the barrier diffusion. We set $\\sqrt{a}=0.89$, which is the value inferred by MR in \\Cite{Maggiore:2009rw} using the simulations of Robertson \\etal~\\cite{Robertson:2008jr}. We wish to emphasize a feature that our calculation shares with that of MR, which is that the constant $a$ is the only parameter whose value depends on the output of $N$-body simulations. The rest of the calculation for the mass function is completely analytical and from first principles.  To make the comparison meaningful, we introduce theoretical error bars on the curves. These error bars have no intrinsic statistical meaning -- they simply keep track of the absolute magnitude of the terms that are ignored in any given prescription for computing the mass function. As we have discussed at length in section \\ref{consistency}, these theoretical errors are scale dependent. The estimated error magnitude for each point is the maximum among the terms ignored in the expression. More explicitly, the errors for the linearized LMSV expression~\\eqref{f-LMSVlin} are estimated as the maximum of $(\\ep\\nu^3)^2$ which comes from the expansion of the exponential, $\\ep\\nu$ which is the order of the largest unequal time terms missing, and $\\tilde{\\kappa}\\nu^{-2}$ which comes from the filter effects. The errors for the LMSV expression~\\eqref{f-LMSVquad} are similarly estimated as the maximum of $(\\ep\\nu^3)^3$, $\\ep\\nu$ and $\\tilde{\\kappa}\\nu^{-2}$. The largest error for the MVJ expression~\\eqref{f-MVJ} is the maximum of $\\ep\\nu$ (unequal time terms) and $\\tilde{\\kappa}\\nu^{-2}$ (filter effects). Finally, the error for the MR expression~\\eqref{f-MR-detailed} is the maximum of $(\\ep\\nu^3)^2$ from the expansion of the exponential, $\\ep \\nu^{-3}$ from the largest unequal time terms ignored, and $\\tilde{\\kappa}^2 \\nu^{-2}$ and $\\tilde{\\kappa} \\ep \\nu$ from the filter effects. We include the filter effects and the associated errors only in~\\fig{final-comparison-z1k}.   \\begin{figure}[t] \\centering \\includegraphics[height=0.35\\textheight]{mf_z1} \\caption{\\footnotesize Theoretical comparison of the different mass functions at $z=1$, without the filter effects, i.e. setting $\\ti{\\kappa}=0$. We plot the ratio $R_{\\rm ng}$ of the non-Gaussian and Gaussian mass functions, in the local model with $\\fnl=100$ and $\\gnl=0$. See main text for a discussion of the error bars. The arrow indicates the mass scale where $\\eps_1 \\nu^3/6 = 1$, i.e. where the expansions of LMSV (both linearized and quadratic) and MR break down.} \\label{final-comparison-z1} \\end{figure}  \\begin{figure}[t] \\centering \\includegraphics[height=0.35\\textheight]{mfk_z1} \\caption{\\footnotesize Same as \\fig{final-comparison-z1}, but including filter effects. These affect only the error bars for MVJ and LMSV, and they affect both the curve and the error bars for MR and our result. For MR and our result, the Gaussian mass function used to construct the ratio $R_{\\rm ng}$, is taken as the non-Gaussian result at $\\fnl=0$, and hence includes filter effects.} \\label{final-comparison-z1k} \\end{figure}  From these figures, we can draw some interesting conclusions. First of all, we see that it is important to retain terms which are quadratic in the NG, either with a saddle point method like in MVJ and in our formula, or by expanding the exponential up to second order, like in LMSV. Actually, we argue that it is correct to keep the exponential, otherwise the expansion breaks down when $\\ep \\nu^3$ is of order unity. We notice in passing that the term proportional to $\\eps_2$ which comes from the trispectrum, partially cancels with the $\\eps_1^2$ term. Secondly, comparing our expression with MVJ's, we can observe that keeping the unequal time terms allows us to sensibly reduce the theoretical errors due to the approximate treatment of NG. In fact, if these terms are missing, they provide the largest theoretical error on large scales. Instead, the largest theoretical error on small scales comes from the approximations involved in dealing with a real space top-hat filter, as is apparent from~\\fig{final-comparison-z1k}.   To conclude, in this work we have calculated the non-Gaussian halo mass function in the excursion set framework, improving over previous calculations. We started from a path integral formulation of the random walk of the smoothed density field, following~\\Cite{Maggiore:2009rv}. This allows us to take into account effects due to multi-scale correlations of the smoothed density field (``unequal time'' correlations), and due to the real space top-hat filter, which generates non-Markovianities in the random walk. We recognize two small parameters in which we perturb: $\\ep$, defined below \\eqn{notes-eq1}, which measures the magnitude of the primordial NG; and $\\nu^{-1} = \\sigma_R/\\del_c$, which is small on very large scales. In order to do a consistent expansion and to estimate the theoretical errors, one must study the (scale dependent) relation between these two parameters, which we have discussed in Sec.~\\ref{consistency}. We then used saddle point techniques which allowed us to non-perturbatively retain the dependence on $\\ep \\nu^3$, which naturally appears in the calculation and whose magnitude becomes of order unity at high masses and high redshift. Finally, we included effects due to the choice of filter function and due to deviations from spherical collapse, as explained in Sec.~\\ref{diffuse-filter}. Our final result is presented in~\\eqn{f-fullcorrected}, which we reproduce here: \\begin{align} f(\\nu,t) = f_{\\rm PS}(\\nu)&\\bigg(1- \\ti{\\kappa} + \\Cal{O}(\\ti{\\kappa}^2) \\bigg) \\exp\\left[\\frac{1}{6}\\eps_1\\nu^3 - \\frac{1}{8}\\left(\\eps_1^2-\\frac{\\eps_2}{3}\\right)\\nu^4 \\right]\\nonumber\\\\ &\\times\\bigg\\{ 1 + \\frac{\\left(1-2d\\ln\\ti{\\kappa}/d\\ln t\\right)}{1-\\ti{\\kappa}}\\ti{\\kappa} \\nu^{-2} \\left( 1 -  2\\nu^{-2}\\right) -\\frac{1}{4}\\eps_1\\nu\\left(\\frac{c_1}{1-\\ti{\\kappa}} +  4-2c_1 \\right) \\nonumber\\\\ &\\ph{\\times\\bigg\\{ 1+} -\\frac{1}{4}\\eps_1\\nu^{-1}\\left(c_1-\\frac{1}{4}c_2-2\\right) \\nonumber\\\\ &\\ph{\\times\\bigg\\{ 1+-1} + \\Cal{O}(\\ti{\\kappa}^2\\nu^{-2}, \\ti{\\kappa}\\ep\\nu,\\ti{\\kappa}\\nu^{-6}) + \\Cal{O}(\\ep^2\\nu^2,\\ep^3\\nu^5,\\ep\\nu^{-3}) \\bigg\\} \\, . \\la{f-fullcorrected-repeat} \\end{align} In \\tab{fits} we provide analytical approximations for the various parameters that appear in this expression, which in general must be computed numerically. We also considered other expressions for the mass function found in the literature, which use different expansion methods but do not estimate the theoretical errors. We estimated the theoretical errors for each formula, and we show comparative plots in \\fig{final-comparison-z1} and \\fig{final-comparison-z1k}. In our work we have improved over the calculations of MVJ \\cite{Matarrese:2000iz} and LMSV \\cite{LoVerde:2007ri} (who ignore unequal time correlations) and of MR \\cite{Maggiore:2009rx} (who do not retain the exponential dependence on $\\ep\\nu^3$). We have also demonstrated that the (linearized) result of LMSV can be significantly improved by retaining the quadratic terms of their calculation which are usually ignored in the literature. We find that at large scales and high redshifts, the biggest theoretical errors are introduced by ignoring the exponential dependence on $\\ep\\nu^3$, followed by the neglect of unequal time correlations. The errors on our expression \\eqref{f-fullcorrected-repeat} are therefore significantly smaller than those of the others. The strength of our approach lies in the combination of path integral methods as laid out by MR \\cite{Maggiore:2009rx}, and the saddle point approximation as used by MVJ \\cite{Matarrese:2000iz}.  Our work can be continued in several directions. First, a thorough calculation of the effects due to the choice of the filter should be performed, since they lead to significant uncertainties in our final expression. This would include a study of the details of the continuum limit of the path integral near the barrier, and also a study of the statistics of the barrier diffusion process in the presence of filter effects. Second, a comparison with $N$-body simulations should be performed, in order to quantitatively assess the possibility of constraining NG using our work. Third, it would be interesting to study how to account for the effects of ellipsoidal collapse, in a framewrok such as the one employed in this paper.  Finally, an application to the void statistics along the same lines should be feasible. The problem here is made more interesting by the presence of \\emph{two} barriers, as discussed by Sheth \\& van de Weygaert \\cite{Sheth:2003py}, and since voids probe larger length/mass scales than halos, they constitute a promising future probe of primordial NG \\cite{Kamionkowski:2008sr}."
    },
    "1005/1005.5312.txt": {
        "abstract": "We present a novel radiation hydrodynamics code, START, which is a smoothed particle hydrodynamics (SPH) scheme coupled with accelerated radiative transfer. The basic idea for the acceleration of radiative transfer is parallel to the tree algorithm that is hitherto used to speed up the gravitational force calculation in an $N$-body system. It is demonstrated that the radiative transfer calculations can be dramatically accelerated, where the computational time is scaled as $N_{\\rm p} \\log N_{\\rm s}$ for $N_{\\rm p}$ SPH particles and $N_{\\rm s}$ radiation sources. Such acceleration allows us to readily include not only numerous sources but also scattering photons, even if the total number of radiation sources is comparable to that of SPH particles. Here, a test simulation is presented for a multiple source problem, where the results with START are compared to those with a radiation SPH code without tree-based acceleration. We find that the results agree well with each other if we set the tolerance parameter as $\\theta_{\\rm crit}\\le1.0$, and then it demonstrates that START can solve radiative transfer faster without reducing the accuracy. One of important applications with START is to solve the transfer of diffuse ionizing photons, where each SPH particle is regarded as an emitter. To illustrate the competence of START, we simulate the shadowing effect by dense clumps around an ionizing source. As a result, it is found that the erosion of shadows by diffuse recombination photons can be solved. Such an effect is of great significance to reveal the cosmic reionization process. ",
        "introduction": "The radiative transfer (RT) in three-dimensional space is virtually a six-dimensional problem for a photon distribution function in the phase space. So far, various RT schemes have been proposed, some of which are coupled with hydrodynamics \\citep{Iliev06,Iliev09}. In a grid-based scheme, the radiative transfer can be reduced to a five-dimensional problem without the significant reduction of accuracy [e.g. the ART scheme that is proposed in \\citet{Nakamoto01} and \\citet{Iliev06}]. In a smoothed particle hydrodynamics (SPH) scheme, if diffuse scattering photons are neglected, the computational cost is proportional to $N_{\\rm p}N_{\\rm s}$, where $N_{\\rm p}$ and $N_{\\rm s}$ are the number of SPH particles and that of radiation sources, respectively. This type of radiation transfer solver can be coupled with the hydrodynamics (e.g. Susa 2006; Illiev et al. 2006, 2009).  But, in various astrophysical problems, diffuse scattering photons play a significant role, for example, in photoionization of a highly clumpy medium. The treatment of diffuse radiation is a hard barrier owing to its high computational cost. Susa (2006) proposed Radiation-SPH ({RSPH}) scheme, in which the radiation transfer based on ray-tracing is coupled with SPH. Since SPH particles are directly used to integrate optical depths, high density regions can be automatically resolved with high accuracy. The code can actually treat multiple radiation sources, but the computational cost increases in proportion to the number of radiation sources. Hence, it is difficult to include numerous radiation sources. For the same reason, recombination photons are hard to involve, since each SPH particle should be treated as a radiation source. Thus, in the applications of RSPH, the on-the-spot approximation is assumed \\citep{Spitzer78}, where recombination photons are supposed to be absorbed on the spot, and therefore transfer of diffuse radiation is not solved. The RSPH has been applied to explore the radiative feedback on first generation object formation \\citep{SU04, SU06, Susa07, Susa08, HUS09, SUH09}, where ionization of hydrogen and photo-dissociation of hydrogen molecules $\\rm H_2$ by ultraviolet (UV) radiation are key physics. But, diffuse radiation potentially changes the ionization fraction near ionization-front (I-front), and therefore alter $\\rm H_2$ abundance. Also, diffuse radiation can erode neutral shadows behind dense clumps. Such erosion is significant to elucidate the cosmic reionization process. To treat the diffuse radiation properly, the acceleration of radiation transfer solver is indispensable.  In this paper, we present a novel radiation transfer solver for an SPH scheme, START, with the tree-based acceleration of the radiative transfer. The paper is organized as follows. In Section 2, we describe the algorithm in detail. In Section 3, we test the new scheme by comparing with a previous scheme, for the propagation of ionization fronts around multiple sources. In Section 4, we demonstrate the impacts of diffuse recombination photons by solving the transfer of diffuse radiation. Section 5 is devoted to the conclusions. Some discussion on the related topics is also given there. ",
        "conclusions": "We have presented a novel accelerated radiation hydrodynamics scheme START, which is an SPH scheme with tree-based acceleration of radiation transfer. In this scheme, an oct-tree structure is utilized to reduce the effective number of radiation sources. The computational time by START is roughly proportional to $N_{\\rm p}\\log N_{\\rm s}$, where $\\log N_{\\rm s}$ is the radiation source number and $N_{\\rm p}$ is the SPH particle number. We have shown the accuracy of START is almost equivalent to that of a previous radiation SPH scheme, if we set the tolerance parameter as $\\theta_{\\rm crit}\\le1.0$. The method presented in this paper provides a powerful tool to solve numerous radiation sources. Also, the new scheme allows us to solve the transfer of diffuse scattering photons. In this case, the computational time is proportional to $N_{\\rm p}\\log N_{\\rm p}$. We have simulated the expansion of an HII region around a radiation source in an initially uniform medium. Comparing these results with 1D spherical symmetric RHD, we have confirmed that our new method correctly solve the transfer of diffuse recombination photons. In addition, we have explored the impacts of recombination photons in the case that dense clumps are irradiated by a luminous radiation source. As a result, we have found that the recombination photons can dramatically change the ionization structure especially behind clumps, if the size of the clump is smaller than the mean free path of the recombination photons.  START is likely to be available for various astrophysical issues in which numerous radiation sources are required. One of them is the simulation of cosmic reionization. Many simulations have hitherto been widely performed [e.g.,  \\cite{Baek09}, \\cite{Iliev07}, \\cite{Mellema06}, \\cite{Santos08}, and \\cite{Thomas09}]. With START, it would be possible to trace the reionization process by solving hydrodynamics consistently coupled with the transfer of UV photons from numerous sources. Furthermore, the diffuse radiation may play an important role when we consider the impacts of UV radiation from a massive star on neighboring medium \\citep{AS07, Whalen08, HUS09, HUK09}. We will further challenge these issues in forth-coming papers."
    },
    "1005/1005.2170_arXiv.txt": {
        "abstract": "We present full evolutionary calculations appropriate for the study of hydrogen-rich DA white  dwarfs.  This is done by  evolving white dwarf progenitors from the zero age main sequence, through the core hydrogen burning  phase, the  helium burning  phase and  the  thermally pulsing asymptotic  giant branch  phase to  the white  dwarf  stage.  Complete evolutionary sequences are computed for a wide range of stellar masses and for two different metallicities: $Z=0.01$, which is representative of the solar neighborhood, and $Z=0.001$, which is appropriate for the study of old stellar systems, like globular clusters. During the white dwarf cooling  stage we compute  self-consistently the phase  in which nuclear reactions are still  important, the diffusive evolution of the elements in the  outer layers and, finally, we  also take into account all  the relevant energy  sources in  the deep  interior of  the white dwarf,  like   the  release  of   latent  heat  and  the   release  of gravitational  energy  due  to  carbon-oxygen  phase  separation  upon crystallization.   We also  provide  colors and  magnitudes for  these sequences, based on  a new set of improved  non-gray white dwarf model atmospheres, which  include the  most up-to-date physical  inputs like the Ly$\\alpha$ quasi-molecular opacity.  The calculations are extended down to an effective temperature  of 2,500 K. Our calculations provide a homogeneous set of  evolutionary cooling tracks appropriate for mass and  age determinations of  old DA  white dwarfs  and for  white dwarf cosmochronology of the different Galactic populations. ",
        "introduction": "\\label{intro}  White dwarfs are the most common stellar evolutionary end-point.  As a matter of  fact, more  than 97\\%  of stars are  expected to  end their lives as  white dwarfs. Their evolution  can be described  as a simple cooling process which lasts for  very long periods of time.  Moreover, since  these fossil  stars are  abundant and  long-lived  objects they convey  important information  about  the properties  of all  galactic populations (Hansen \\& Liebert  2003). In particular, white dwarfs can be considered  as reliable cosmic  clocks to infer  the age of  a wide variety of  stellar populations,  such as the  Galactic disk  and halo (Winget  et al.   1987; Garc\\'\\i  a--Berro et  al.  1988a,b;  Isern et al. 1998; Garc\\'\\i a--Berro et al. 1999; Torres et al.  2002), and the system of globular  and open clusters (Kalirai et  al. 2001; Hansen et al. 2002, 2007).  They also provide helpful information about the past history of the  star formation rate of our  Galaxy (D\\'\\i az--Pinto et al.  1994).   Additionally, they can  be used to place  constraints on exotic elementary particles (Isern et al. 1992; C\\'orsico et al. 2001; Isern et al. 2008) or on alternative theories of gravitation (Garc\\'\\i a--Berro et al.  1995; Benvenuto  et al. 2004).  Recent reviews on the properties and evolution of white dwarfs and of their applications are those  of Winget  \\& Kepler  (2008) and  Fontaine \\&  Brassard (2008). However, to use them for  this wide range of applications reliable and accurate evolutionary sequences must be computed.  Previous  recent  evolutionary  calculations  of  hydrogen-rich  white dwarfs  are those of  Fontaine et  al. (2001),  Salaris et  al. (1997; 2000),  Hansen (1998;  1999),  Benvenuto \\&  Althaus  (1999) and  Wood (1992; 1995).  All  these works have studied different  aspects of the evolution of white dwarfs with hydrogen-rich envelopes.  For instance, the  most commonly used  models, those  of Wood  (1995), cover  a wide range of stellar masses and  envelope masses and, until recently, were considered  to be a  standard reference  in the  field of  white dwarf cosmochronology.   However,  we   emphasize  that  these  models  were computed using gray atmospheres,  a severe drawback (especially at low luminosities) that more recent calculations have overcome. Among these models  we mention  the works  of Hansen  (1998; 1999).   This  set of cooling models  pioneered the usage  of detailed model  atmospheres as surface  boundary conditions  in  cooling calculations  for old  white dwarfs. This  is an important issue  since it affects  the location of the base  of the convective  envelope. These calculations  also showed that collision-induced  absorption processes affect the  colors of old white dwarfs. Salaris  et al. (1997; 2000) focused  on the question of the interior  abundances of  carbon and oxygen,  which is  of critical importance  to  derive reliable  ages  and  on  the effects  of  phase separation of  carbon and oxygen upon  crystallization, while Fontaine et al.  (2001) were the first to discover the importance of convective coupling between  the atmosphere and  the degenerate core.   This last issue  also  bears  importance   for  the  determination  of  accurate ages. Nevertheless, all these  works suffer from the same shortcoming. All of them  evolved intial white dwarf configurations  which were not obtained   self-consistently    from   models   evolving    from   the main-sequence.  As  a consequence, the chemical  stratification of the inner, degenerate core was simplistic  in most of the cases, except in the case  of the  cooling sequences of  Salaris et al.   (1997; 2000). Also, the envelope mass and the outer layer chemical distribution were idealized ones  in all the cases.  Additionally,  the oldest sequences used physical inputs which are nowadays outdated and, finally, most of them disregarded the energy  release of carbon-oxygen phase separation (Garc\\'\\i a--Berro et al. 1988a; 1988b).  The aim of this paper is to compute a set of new cooling sequences for hydrogen-rich white dwarfs, incorporating the most up-to-date physical inputs.  We emphasize that our evolutionary sequences are derived from a  full and  self-consistent  treatment of  the complete  evolutionary history  of progenitor stars  with different  masses evolved  with two different metallicities  ($Z=0.01$ and $Z=0.001$)  appropriate for the study  of the solar  neighborhood and  of metal-poor  stellar systems, like globular  clusters or the galactic halo.   Thus, our calculations constitute a  comprehensive set of evolutionary tracks  which allow to study the evolution of hydrogen-rich white dwarfs in a wide variety of stellar  systems.  Moreover,  since  our  calculations  encompass  the pre-white  dwarf  evolutionary phases,  the  white dwarf  evolutionary calculations  presented  here  are  not  affected  by  inconsistencies arising from  artificial procedures  to generate starting  white dwarf configurations.   In particular, the  calculation of  the evolutionary history of progenitor  stars provides us with the amount  of H left in the white dwarf, and with  the chemical profiles expected not only for the carbon-oxygen core, but  also for the partially degenerate regions above the core, of relevance for the white dwarf cooling phase.  The paper  is organized as follows.  In  Sect.  \\ref{Computational} we describe the main physical inputs of our models.  We also describe the main  characteristics of  the  initial models,  the model  atmospheres employed  in   this  work  and   some  details  of   the  evolutionary computations.  In  Sect.  3  we present the  results of  the pre-white dwarf evolutionary sequences, paying  special attention to the derived initial-to-final-mass relationship,  and also the  white dwarf cooling tracks. In this section we also discuss in detail the effects of phase separation upon crystallization.  Finally, in Sect. 4 we summarize our results, we discuss their significance and we draw our conclusions. ",
        "conclusions": "\\label{conclusiones}  We have  computed a set  of cooling sequences for  hydrogen-rich white dwarfs,   which   are    appropriate   for   precision   white   dwarf cosmochronology.     Our     evolutionary    sequences    have    been self-consistently evolved from the zero age main sequence, through the core hydrogen and helium  burning evolutionary phases to the thermally pulsing asymptotic  giant branch and,  ultimately, to the  white dwarf stage.   This has  been  done  for white  dwarf  progenitors with  two different metallicities.  For the  first set of evolutionary sequences we have adopted solar metallicity.   This allows us to obtain accurate ages for  white dwarfs in the  local Galactic disk. The  second set of cooling  sequences  corresponds  to  a  metallicity  typical  of  most Galactic  globular  clusters,   $Z=0.001$,  thus  allowing  to  obtain accurate  ages for  metal-poor stellar  systems.  To the  best of  our knowledge,  this  is the  first  set  of self-consistent  evolutionary sequences covering different initial masses and metallicities.  Our  main   findings  can  be  summarized  as   follows.   First,  our evolutionary    sequences    correctly    reproduce    the    observed initial-to-final mass relationship  of white dwarfs. Specifically, our evolutionary calculations  are in excellent agreement  with the recent results  of  Salaris  et  al.   (2009) for  white  dwarfs  with  solar metallicity  progenitors.  Second, we  corroborate  the importance  of convective coupling at low luminosity  in the cooling of white dwarfs, as  originally  suggested  by  Fontaine  et  al.  (2001).   Third,  we demonstrate  the  importance of  residual  hydrogen  burning in  white dwarfs resulting from  low-metallicity progenitors. Fourth, we confirm as  well  the  importance   of  carbon-oxygen  phase  separation  upon crystallization,  in good  qualitative agreement  with the  results of Garc\\'\\i a--Berro et al.  (1988a;  1988b), Segretain et al. (1994) and Salaris et al. (1997; 2000).  Although the computed delays are smaller than those previously estimated by  Segretain et al.  (1994), they are larger than  those computed by Salaris  et al.  (2000), and  are by no means  negligible if precision  white dwarf  cosmochronology is  to be done.   However, we  would like  to mention  that these  delays depend crucially  on  the  previous   evolutionary  history  of  white  dwarf progenitors    and,    particularly,    on    the    rate    of    the $^{12}$C$(\\alpha,\\gamma)^{16}$O  nuclear reaction, as  well as  on the adopted  treatment  for convective  mixing.   Additionally, since  our evolutionary  sequences   rely  on  state-of-the-art   non-gray  model atmospheres,  they reproduce  the  well-known blue  hook  of very  old hydrogen-rich  white dwarfs  caused  by H$_2$-H$_2$  collision-induced absorption  (Hansen et  al.  1999).  Finally, we  show  the impact  of Ly$\\alpha$  quasi-molecular opacity  on  the evolution  of cool  white dwarfs in the color-magnitude diagram.  We  would like  to emphasize  that our  full treatment  of  the entire evolutionary  history  of  white  dwarfs  has  allowed  us  to  obtain consistent white dwarf initial configurations. In particular, the mass of the hydrogen-rich  envelope and of the helium  buffer were obtained from evolutionary  calculations, instead  of using typical  values and artificial initial white dwarf  models.  This has implications for the cooling rates of  old white dwarfs, as the  thicknesses of these outer layers  control  the cooling  speed  of  such  white dwarfs.   Another important  issue  which   we  would  like  to  mention   is  that  our calculations  also yield  self-consistent interior  chemical profiles. This  also has  relevance  for the  cooling  of white  dwarfs, as  the release of  latent heat and gravitational energy  due to carbon-oxygen phase separation upon crystallization crucially depend on the previous evolutionary   history   of   white   dwarfs.   Also,   the   chemical stratification of  white dwarf progenitors is  important for correctly computing the specific  heat of white dwarf interiors.   We would like to stress as well that the evolutionary tracks of cooling white dwarfs presented  here has  been  computed with  the  most accurate  physical inputs and  with a high degree  of detail and  realism. In particular, our calculations include  nuclear burning at the very  early phases of white  dwarf evolution  --- which  is important  to determine  the final thickness   of   the   hydrogen-rich   envelope  ---   diffusion   and gravitational settling  --- which are important to  shape the profiles of the  outer layers  --- accurate neutrino  emission rates  --- which control the cooling at high luminosities --- crystallization and phase separation of carbon  and oxygen --- which dominate  the cooling times at low luminosities --- a very detailed equation of state --- which is important  in all the  evolutionary phases  --- and  improved non-gray model atmospheres --- which allow for a precise determination of white dwarf colors  and outer boundary  conditions for the  evolving models. Finally, we would like to remark that these evolutionary sequences are important  as well for  the calculation  of self-consistent  models of pulsating DA  white dwarfs. All in  all, we consider  that our cooling sequences  represent a  landmark  in the  calculation  of white  dwarf evolutionary tracks, and they open an new era to precision white dwarf cosmochronology.  Detailed tabulations of our evolutionary sequences are available at our Web site {\\tt http://www.fcaglp.unlp.edu.ar/evolgroup}."
    },
    "1005/1005.2352_arXiv.txt": {
        "abstract": "It is discovered in our previous work that different observational systematics, e.g., errors of antenna pointing directions, asynchronous between the attitude and science data, can generate pseudo-dipole signal in full-sky maps of the cosmic microwave background (CMB) anisotropy published by The Wilkinson Microwave Anisotropy Probe (WMAP) team. Now the antenna sidelobe response to the Doppler signal is found to be able to produce similar effect as well. In this work, independent to the sources, we uniformly model the pseudo-dipole signal and remove it from published WMAP7 CMB maps by model fitting. The result demonstrates that most of the released WMAP CMB quadrupole is artificial. ",
        "introduction": "The WMAP mission makes measurements of temperature with two antennas, A and B, and record in time-order the difference between the two antenna temperatures, $T_A-T_B$, which is called the time-order data (TOD). The observed CMB signal is contaminated by Doppler effect induced by the joint motion of the solar system and the spacecraft. The aroused dipole difference signal can be calculated by \\begin{equation}\\label{equ:dp1} d=\\frac{T_0}{c}~\\bf v\\cdot (\\bf n_{_A}-\\bf n_{_B})\\,, \\end{equation} where $T_0=2.725$\\,K is the CMB monopole, $c$ is the speed of light, $\\bf{v}$ is the joint velocity, $\\bf{n}_{_A}$ and $\\bf{n}_{_B}$ are the unit direction vectors of the antenna A and B respectively ~\\citep{hin09}. In this paper the vectors $\\bf v$, $\\bf n_{_A}$, and $\\bf n_{_B}$ are defined in the spacecraft coordinate system. The errors of line-of-sight (LOS) vector, $\\Delta \\bf{n}_{_A}$ and $\\Delta \\bf{n}_{_B}$, will produce a pseudo-dipole difference signal \\begin{eqnarray}\\label{equ:pds} d^* &=& \\frac{T_0}{c}~\\bf{v}\\cdot (\\rm\\Delta\\bf{n}_{_A}-\\rm\\Delta\\bf{n}_{_B}) \\nonumber \\\\ &=&\\frac{T_0}{c}~\\bf{v}\\cdot \\rm\\Delta\\bf{n}\\,. \\end{eqnarray}  The dipoles for each observations have to be removed from the raw data before map-making because their intensities are roughly 10 to 20 times greater than those of the CMB anisotropies. A small error of antenna direction will produce an error in predicted dipole intensity and then cause a pseudo-dipole signal in the resulting CMB map more noticeably through the Doppler dipole subtraction process\\footnote{A true and pure time-order dipole signal produces only a dipole component in the final CMB map, but a time-order pseudo-dipole signal can produce more components, especially quadrupole, on the final CMB map. See Fig.~1.} .  For example, a LOS error of $\\sim7'$, just about a half-pixel in the WMAP resolution, can consequently cause the dipole signal to be deviated by 10--20\\,$\\mu$K, which can not be ignored compared to the very weak CMB signal. An asynchronous between the attitude and differential data can also produce the pseudo-dipole signal. The WMAP mission uses two separate clocks for the attitude data and science data respectively. Therefore, if there is a small constant timing error, there will be a constant direction difference between the \"observed pixel\" and the true pixel. This has the same effect as a constant LOS error in spacecraft coordinates \\citep{liu10}. ",
        "conclusions": "We have found that at least three systematical effects inducing pseudo-dipole signal in WMAP raw differential temperature data can be summarized into a uniform description: the equivalent differential direction error $\\Delta\\bf{n}$. There could be more sources with similar attributes to these three, but they can all be covered equivalently by a simple $\\Delta\\bf{n}$. For example, the LOS is determined by observing the Jupiter; however, for each observation, the relative direction of Jupiter provides no more than three conditions: $(x,y,z)$, but we need four conditions to definitely determine the attitude of the spacecraft: $(x,y,z)$ and rotation. What's more, the LOS is determined together with the antenna main beam response; therefore, the uncertainty of the beam response is also mixed into the LOS uncertainty. Moreover, even the WMAP team themselves have found that some systematical effect can produce quadrupole-like deviations on the final CMB map (Jarosik et al. 2007, Fig. 3). Recently, \\citet{rou10} pointed out that a small timing error during the step of calibration of the raw TOD could correspond to adding a pseudo-dipole difference signal with no effect on positional data. It is impossible to exactly estimate all such effects, especially when some of them could be unknown for now, but it is easy to model the impact of $\\Delta\\bf{n}$ upon the final CMB map and remove it by model fitting. This is apparently much more feasible and effective than estimating and removing all possible sources one by one.  What is important for model fitting is the feature of a template map, not the absolute amplitude. For example, the WMAP foreground removal will be the same even if the three foreground templates $t_d$, $t_{_H}$ and  $t_s$ are all doubled. In this work, the result of removing pseudo-dipole signal is not dependent to what amplitude is assumed for the direction error: larger amplitude leads to lower fitted coefficient in model fitting and the final correction to the CMB map is the same. Here it is worth to mention that the characteristic feature of template map of the pseudo-dipole signal induced by a pointing or timing error obtained from our previous work \\citep{liu10} or shown in Fig.~1 of this work  has already been  confirmed by other independent works \\citep{mos10,rou10}.  Our result demonstrates  that a pseudo-dipole field can produce deviations in the sky temperature map with a structure very similar to the CMB quadrupole pattern published by the WMAP team. The template map of pseudo-dipole induced temperature deviation shown in Fig.~1 are generated from a differential dipole field which is completely determined by the spacecraft velocity and overall-equivalent error of differential direction without using any CMB signal.  The pseudo-dipole signal in the differential datum for a WMAP observation distorts the temperatures for corresponding sky pixels varying observation by observation.  Therefore, the pseudo-dipole induced temperature map is generated by the pseudo-dipole signal combined with the WMAP scan pattern,  which is highly relative to the ecliptic plane~\\citep{hin07, li09}.   The WMAP CMB quadrupole component being highly aligned and close to the ecliptic plane is a long time puzzle in cosmology, which now can be naturally explained by the pseudo-dipole signal effect.  Keeping insist that the published WMAP quadrupole is really cosmological origin now becomes more difficult:  it is needed the primordial density fluctuations not only occasionally being laid down in the plane of the solar system, but also occasionally having almost the same phase as the scan strategy of the WMAP mission!"
    },
    "1005/1005.4604_arXiv.txt": {
        "abstract": "Starting in winter 2008/2009 an \\textit{L}-band 7-Feed-Array receiver is used for a 21-cm line survey performed with the 100-m telescope, the Effelsberg--Bonn \\ion{H}{i} survey (EBHIS). The EBHIS will cover the whole northern hemisphere for decl.$ > -5\\degr$ comprising both the galactic and extragalactic sky out to a distance of about 230\\,Mpc. Using state-of-the-art FPGA-based digital fast Fourier transform spectrometers, superior in dynamic range and temporal resolution to conventional correlators, allows us to apply sophisticated radio frequency interference (RFI) mitigation schemes.  In this paper, the EBHIS data reduction package and first results are presented. The reduction software consists of RFI detection schemes, flux and gain-curve calibration, stray-radiation removal, baseline fitting, and finally the gridding to produce data cubes. The whole software chain is successfully tested using multi-feed data toward many smaller test fields (1--100 deg$^2$) and recently applied for the first time to data of two large sky areas, each covering about 2000 deg$^2$. The first large area is toward the northern galactic pole and the second one toward the northern tip of the Magellanic Leading Arm. Here, we demonstrate the data quality of EBHIS Milky Way data and give a first impression on the first data release in 2011. ",
        "introduction": "Blind \\ion{H}{i} surveys allow us to explore a variety of objects. The Effelsberg--Bonn \\ion{H}{i} Survey (EBHIS) covers not only the Milky Way but also the local universe out to a redshift of 0.07. Accordingly, EBHIS will serve as a major database for many scientific questions, addressing, e.g., the structure and mass distribution of the Milky Way, the size distribution of halo clouds, or the structure formation of the local universe. A complete description of the scientific aims will be given in a forthcoming paper. Here, the technical setup and the data reduction software as used for the survey are presented.  Today, the most comprehensive database for galactic \\ion{H}{i} science is the Leiden/Argentine/Bonn survey \\citep[LAB;][]{kalberla05} the first all-sky survey corrected for stray radiation. Toward high galactic latitudes, the faint Milky Way emission from the direction of interest can be severely degraded by radiation of the Milky Way disk entering the receiver system via the near  and far side lobes. SR correction is crucial for quantitative analyses of most of the galactic sky using conventional single dish optics. The Parkes and Arecibo telescopes perform multiple consecutive surveys to observe the galactic and extragalactic sky. The data acquisition of the Parkes Galactic All Sky Survey \\citep[GASS;][]{mcclure06,mcclure09} is already completed and the final data products are released \\citep{kalberla10}. The Galactic ALFA \\citep[GALFA;][]{goldsmith04,heiles04} is still in progress. In the extragalactic regime the \\ion{H}{i} Parkes All Sky Survey \\citep[HIPASS;][]{barnes01,meyer04} mapped the complete southern hemisphere detecting more than 5000 galaxies providing a valuable database to infer the \\ion{H}{i} properties of galaxies in the local universe. The Arecibo Legacy Fast ALFA survey \\citep[ALFALFA;][]{giovanelli05p} is ongoing and will result in deeper data at higher angular resolution though it is limited to a smaller area of about 7\\,000 deg$^2$.    Since the beginning of the year 2009 an \\textit{L}-band 7-Feed-Array receiver system is available at the 100-m telescope in Effelsberg for astronomical measurements. We use this instrument to perform a fully sampled \\ion{H}{i} survey of the northern hemisphere for $\\mathrm{decl.}> -5\\degr$, observing in parallel the galactic and extragalactic sky. Multiple coverages (at different hour angles and seasons) will allow us to disentangle radio frequency interference (RFI), SR, and baseline problems. Redundancy allows us to remove instrumental biases significantly. The survey area is subdivided into two sky areas, one toward the Sloan Digital Sky Survey \\citep[SDSS;][]{adelman08} area where an integration time of 10\\,minutes is projected while the remaining sky will be integrated for 2\\,minutes. This yields a full-sky survey superior in sensitivity, angular sampling, and resolution to any previously performed large \\ion{H}{i} survey \\citep{kalberla09}. The data of the first full sky coverage will be released in 2011.   The data analysis of a large-area \\ion{H}{i} survey is undoubtedly a major challenge. Modern receiving systems offer scientifically usable bandwidths of several tens to hundreds of MHz to study the \\ion{H}{i} at moderate red shifts. Digital high-dynamic-range spectrometers based on Field Programmable Gate Arrays \\citep[FPGAs; ][]{stanko05,klein06} allow us to store spectra on time scales of less than 1\\,s which is essential to apply sophisticated RFI mitigation procedures \\citep{winkel07a}. The receiving system as used for EBHIS provides a bandwidth of 100\\,MHz which covers the redshift range out to $z\\sim0.07$ while the 16384 spectral channels yield a velocity resolution of $1.25\\,\\mathrm{km\\,s}^{-1}$. Spectral dumps are recorded every 500\\,ms. As a drawback though, one has to deal with very high data rates. For the EBHIS, typical values are about 5\\,GB\\,hr$^{-1}$, making the calibration and analysis of EBHIS data a task which is impossible to accomplish manually. Consequently, the data reduction algorithms and software developed were optimized with respect to computational efficiency and performance.    The paper is organized as follows. In Section\\,\\ref{secinstrument}, the receiving system and technical setup are described. During various test measurements the quality of the new receiving system was investigated, the results of which are discussed in Section\\,\\ref{secreceiverquality}. Our data reduction package which was developed for the EBHIS is presented in detail in Section\\,\\ref{secdatareduction} and we show its successful application to an example data set in  Section\\,\\ref{secfirstresults}. A summary is given in  Section\\,\\ref{secsummary}. ",
        "conclusions": "\\label{secsummary} We presented in this work the current status of the EBHIS data reduction procedures. The observed spectral data contain lots of RFI and exhibit a rather complicated gain curve shape. A lot of effort went into the investigation of these issues and the development of robust and computationally fast algorithms. We discussed the standing wave problem and its potential impact on the widely used frequency or position switching techniques.  The receiver quality was extensively tested. System temperatures are as expected and the complete receiving system has stability times sufficiently high to reach all scientific aims of the EBHIS. Using the standard calibrators S\\,7 or S\\,8, a gain calibration accuracy of better than 3\\% is easily obtained.  The current quality of the EBHIS data and the reduction software is promising. We are confident that even the first data release in 2011 will provide a database in quality comparable to HIPASS and GASS. This will also allow us to produce a great successor to the LAB survey data by combining EBHIS and GASS both having much higher sensitivity and resolution."
    },
    "1005/1005.1217_arXiv.txt": {
        "abstract": "We investigate convergence of Lagrangian Perturbation Theory (LPT) by analysing the model problem of a spherical homogeneous top-hat in an Einstein-deSitter background cosmology.  We derive the formal structure of the LPT series expansion, working to arbitrary order in the initial perturbation amplitude.  The factors that regulate LPT convergence are identified by studying the exact, analytic solution expanded according to this formal structure. The key methodology is to complexify the exact solution, demonstrate that it is analytic and apply well-known convergence criteria for power series expansions of analytic functions.  The ``radius of convergence'' and the ``time of validity'' for the LPT expansion are of great practical interest.  The former describes the range of initial perturbation amplitudes which converge over some fixed, future time interval.  The latter describes the extent in time for convergence of a given initial amplitude. We determine the radius of convergence and time of validity for a full sampling of initial density and velocity perturbations.  This analysis fully explains the previously reported observation that LPT fails to predict the evolution of an underdense, open region beyond a certain time. It also implies the existence of other examples, including overdense, closed regions, for which LPT predictions should also fail. We show that this is indeed the case by numerically computing the LPT expansion in these problematic cases.  The formal limitations to the validity of LPT expansion are considerably more complicated than simply the first occurrence of orbit crossings as is often assumed. Evolution to a future time generically requires re-expanding the solution in overlapping domains that ultimately link the initial and final times, each domain subject to its own time of validity criterion. We demonstrate that it is possible to handle all the problematic cases by taking multiple steps (LPT re-expansion).  A relatively small number ($\\sim 10$) of re-expansion steps suffices to satisfy the time of validity constraints for calculating the evolution of a non-collapsed, recombination-era perturbation up to the current epoch. If it were possible to work to infinite Lagrangian order then the result would be exact.  Instead, a finite expansion has finite errors.  We characterise how the leading order numerical error for a solution generated by LPT re-expansion varies with the choice of Lagrangian order and of time step size. Convergence occurs when the Lagrangian order increases and/or the time step size decreases in a simple, well-defined manner. We develop a recipe for time step control for LPT re-expansion based on these results. ",
        "introduction": "\\label{intro} Understanding the non-linear growth of structure in an expanding universe has been an active area of research for nearly four decades. Simulations have been instrumental in illustrating exactly what happens to an initial power spectrum of small fluctuations but analytic methods remain essential for elucidating the physical basis of the numerical results. Perturbation theory, in particular, is an invaluable tool for achieving a sophisticated understanding.  The Eulerian and Lagrangian frameworks are the two principal modes of description of a fluid. The fundamental dependent variables in the Eulerian treatment are the density $\\rho({\\bf x},t)$ and velocity ${\\bf v}({\\bf x}, t)$ expressed as functions of the grid coordinates ${\\bf x}$ and time $t$, the independent variables.  In perturbation theory the dependent functions are expanded in powers of a small parameter. For cosmology that parameter typically encodes a characteristic small spatial variation of density and/or velocity with respect to a homogeneous cosmology at the initial time.  As a practical matter, the first-order perturbation theory becomes inaccurate when the perturbation grows to order unity.  Subsequently one must work to higher order to handle the development of non-linearity (see \\citealt{bernardeau_large-scale_2002} for a review) or adopt an alternative method of expansion.  In the Lagrangian framework, the fundamental dependent variable is the physical position of a fluid element or particle (terms used interchangeably here).  The independent variables are a set of labels ${\\bf X}$, each of which follows a fluid element, and the time. Usually ${\\bf X}$ is taken as the position of the element at some initial time but other choices are possible. In any case, the physical position and velocity of a fluid element are ${\\bf r} = {\\bf r} ({\\bf X},t)$ and ${\\dot {\\bf r}}({\\bf X},t)$, respectively. Knowledge of the motion of each fluid element permits the full reconstruction of the Eulerian density and velocity fields.  In cosmological applications of Lagrangian perturbation theory (LPT), just like Eulerian perturbation theory, the dependent variables are expanded in terms of initial deviations with respect to a homogeneous background.  The crucial difference is that the basis for the expansion is the variation in the initial position and position-derivative not the variation in the initial fluid density and velocity. The Eulerian density and velocity may be reconstructed from knowledge of the Lagrangian position using exact non-perturbative definitions.  A linear approximation to the displacement field results in a non-linear expression for the density contrast. The Lagrangian description is well-suited to smooth, well-ordered initial conditions; a single fluid treatment breaks down once particle crossings begin, caustics form and the density formally diverges.  First-order LPT was originally introduced by \\cite{zeldovich_gravitational_1970} to study the formation of non-linear structure in cosmology. In his treatment the initial density field was taken to be linearly proportional to the initial displacement field (the ``Zeldovich approximation''). These results were extended by many authors (\\citealt{moutarde_precollapse_1991}; \\citealt{buchert_lagrangian_1992}; \\citealt{bouchet_weakly_1992}; \\citealt{buchert_lagrangian_1993}; \\citealt{buchert_lagrangian_1994}; \\citealt{munshi_non-linear_1994}; \\citealt{catelan_lagrangian_1994}; \\citealt{buchert_lagrangian_1995}; \\citealt{bouchet_perturbative_1995}; \\citealt{bouchet_introductory_1996}; \\citealt{ehlers_newtonian_1997}). The work pioneered by Bouchet focused on Zeldovich initial conditions and established the link between LPT variables and statistical observables.  The work by Buchert as well as the paper by \\cite{ehlers_newtonian_1997} formalised the structure of the Newtonian perturbative series for arbitrary initial conditions.  A general relativistic version of the Zeldovich approximation was developed by \\cite{kasai_tetrad-based_1995} and other relativistic descriptions of the fluid in its rest frame were investigated by \\cite{matarrese_post-newtonian_1996} and \\cite*{matarrese_general-relativistic_1993, matarrese_relativistic_1994}. LPT has been used for many applications including, recently, the construction of non-linear halo mass functions by \\cite{monaco_lagrangian_1997} and \\cite{scoccimarro_pthalos:fast_2002}.  Not much has been written about the convergence of LPT although LPT expansions are routinely employed. \\cite{sahni_accuracy_1996} pointed out that the formal series solution for the simplest problem, the spherical top-hat, did not converge for the evolution of homogeneous voids. \\capfigref{openplot} illustrates the conundrum that the LPT approximations diverge from the exact solution in a manner that worsens as the order of the approximation increases. The details will be described in the next section.  This paper explores LPT convergence for the spherical top-hat and identifies the root cause for the lack of convergence. The analysis naturally suggests a means of extending the range of validity of LPT. This generalisation of LPT guarantees convergence to the exact solution of the model problem at all times prior to the occurrence of the first caustic.   \\cite{tatekawa_improvinglagrangian_2007} attempted to treat the divergence by applying the Shanks transformation to the LPT series. Although non-linear transformations can sum a divergent series, the correct answer is not guaranteed; comparison of several different methods is usually necessary to yield trustworthy results. Other approaches include the Shifted-Time-Approximation (STA) and Frozen-Time-Approximation (FTA) which have been investigated by \\cite*{karakatsanis_temporal_1997}. These schemes modify lower order terms to mimic the behavior of higher order terms and/or extend the range of applicability in time. None of these techniques are considered here.  \\begin{figure} \\includegraphics[width=7cm,height=7cm]{openplot.pdf} \\caption{The time-dependent scale factor $b$ of an initial spherical top-hat perturbation is plotted as a function of the background scale factor $a$. The perturbation is a pure growing mode, i.e. the density and velocity perturbations vanish at $t=0$.  The black dotted line is the exact solution.  The smooth blue lines are the LPT results obtained by  working successively to higher and higher order. Series with even (odd) final order lie below (above) the exact solution.  Roughly speaking, LPT converges only for $a \\lta 0.2$. Beyond that point {\\it the higher order approximations deviate from the exact solution more than lower order ones}.} \\label{openplot} \\end{figure}   The organisation follows: \\S\\ref{LPTforSTH} sketches the model problem, the evolution of a uniform sphere in a background homogeneous Einstein-deSitter cosmology. The LPT equations, the structure of the formal series and the term-by-term solution are outlined. \\S\\ref{convergencestudy} discusses the complexification of the LPT solution and convergence of the series. This section introduces the ``radius of convergence\" and the ``time of validity\" for LPT. \\S\\ref{exactsoln} outlines the real and complex forms of the parametric solution and sets forth the equations that must be solved to locate the poles which govern the convergence. \\S\\ref{Results} presents numerical results for the time of validity and radius of convergence for a full range of possible initial conditions for the top-hat. The notion of mirror model symmetry is introduced and used to explain a connection in the convergence for open and closed models. \\S\\ref{sec_extension} shows that the time of validity may be extended by re-expanding the solution in overlapping domains that ultimately link the initial and final times, each domain subject to an individual time of validity criterion. The feasibility of this method is demonstrated in some examples. \\S\\ref{conclusion} summarises the work. ",
        "conclusions": "\\label{conclusion}  We have investigated the time of validity of Lagrangian perturbation theory for spherical top-hat cosmologies with general initial conditions.  Using techniques from complex analysis we showed that the time of validity is always limited for open models. We also discovered a class of closed models whose time of validity is less than their time to collapse. We introduced the concept of the mirror model and derived a symmetry principle for the time of validity of mirror models.  For small initial perturbations the time of validity of LPT series expansion of an open model corresponds to the collapse time of a closed mirror model.  A qualitative analogy is useful.  A single LPT series expansion is similar to a single step in a finite difference approximation for advancing a hyperbolic partial differential equation like the wave equation. The time of validity of the LPT expansion is analogous to the Courant condition which guarantees stability. In LPT the constraint is an acceleration-related time-scale; in the wave equation it is a sound-crossing time-scale.  We developed the method of LPT re-expansion which overcomes the limitations intrinsic to a single expansion. We demonstrated how to iteratively re-expand the solution so as to link convergent series expressions that extend from initial to final times. The time of validity of the expansions set the minimum number of re-expansion steps ($\\sim 10$) necessary for cosmological simulations starting at recombination and proceeding to the present epoch. Finite as opposed to infinite order Lagrangian expansions required extra steps to achieve given error bounds.  We characterised how the leading order numerical error for a solution generated by LPT re-expansion varied with the choice of Lagrangian order and of time step size. We provided a recipe for time step control for LPT re-expansion based on these results.  Our long-term goal and motivation for this study is to develop a numerical implementation of LPT re-expansion for fully inhomogeneous cosmological simulation. Top-hats with Zeldovich initial conditions have special properties with respect to LPT convergence. We found that all underdense models must be treated by re-expansion while none of the overdense ones need be. However, during the course of an inhomogeneous simulation the density and irrotational velocity perturbations (with respect to a homogeneous background cosmology) at an arbitrary point will generally not fall on the top-hat's Zeldovich curve. Hence, the convergence of LPT in inhomogeneous applications must be guided by the analysis of more general models. Top-hats with arbitrary initial conditions are the simplest possibility and constitute the main focus in this paper. The limitations on LPT convergence which we have elucidated in this generic case are considerably more complicated  than in the top-hat with Zeldovich initial conditions. Our plan is to use the generic time of validity criterion to determine the time-stepping for inhomogeneous evolution. This should allow us to develop high-precision simulations with well-defined control of errors. The practical impact of a refined treatment of LPT convergence is not yet clear.  The convergence issues we have dealt with should not be confused with the breakdown when orbit crossing takes place and the Jacobian of the transformation from Lagrangian to physical coordinates becomes singular. At that time the flow becomes multi-streamed and much of the simplicity and advantage of the Lagrangian approach vanishes.  The aim of the current work is to make sure it is possible to reach the epoch of multi-streamed flow but offers nothing new on how to proceed beyond it. In fact, it may be necessary to include an effective pressure term in the equations to account for the velocity dispersion induced by orbit crossing (\\citealt{adler_lagrangian_1999}; \\citealt*{buchert_extendingscope_1999}) or to adopt alternative approximations for the basic dynamics (such as the adhesion approximation; see \\citealt{sahni_approximation_1995} for a review and references therein) to make progress."
    },
    "1005/1005.4574_arXiv.txt": {
        "abstract": "We explore the cosmological evolution in the exponential gravity $f(R)=R +c_1 \\left(1-e^{- c_2 R} \\right)$ ($c_{1, 2} = \\mathrm{constant}$). We summarize various viability conditions and explicitly demonstrate that the late-time cosmic acceleration following the matter-dominated stage can be realized. We also study the equation of state for dark energy and confirm that the crossing of the phantom divide from the phantom phase to the non-phantom (quintessence) one can occur. Furthermore, we illustrate that the cosmological horizon entropy globally increases with time. ",
        "introduction": "There exist two representative approaches to account for the current accelerated expansion~\\cite{WMAP, Komatsu:2010fb, SN1} of the universe~\\cite{Copeland:2006wr, Nojiri:2006ri, Sotiriou:2008rp, DeFelice:2010aj, Review-Tsujikawa, Linder:2010qn}. One is to introduce ``dark energy'' in the framework of general relativity. The other is to consider a modified gravitational theory, such as $f(R)$ gravity.  Several viable theories of $f(R)$ gravity have been constructed; e.g., power-law~\\cite{Amendola:2006we, Li:2007xn}, Nojiri-Odintsov~\\cite{Nojiri-Odintsov}, Hu-Sawicki~\\cite{Hu:2007nk}, Starobinsky's~\\cite{Starobinsky:2007hu}, Appleby-Battye~\\cite{Appleby:2007vb}, and Tsujikawa's~\\cite{Tsujikawa:2007xu} models (for more detailed references, see a recent review on $f(R)$ gravity~\\cite{DeFelice:2010aj}). It is known that these models can satisfy the following conditions for the viability: (i) positivity of the effective gravitational coupling, (ii) stability of cosmological perturbations~\\cite{Nojiri:2003ft, Dolgov:2003px, Faraoni:2006sy, Song:2006ej}, (iii) asymptotic behavior to the standard $\\Lambda$-Cold-Dark-Matter ($\\Lambda\\mathrm{CDM}$) model in the large curvature regime, (iv) stability of the late-time de Sitter point~\\cite{Amendola:2006we, Muller:1987hp, Faraoni:2005vk}, (v) constraints from the equivalence principle, and (vi) solar-system constraints~\\cite{Solar-System-Constraints}.  Recently, an interesting model of $f(R)=R +c_1 \\left(1-e^{- c_2 R} \\right)$, called ``exponential gravity'', has been proposed in Refs.~\\cite{Exponential-type-f(R)-gravity, Cognola:2007zu, Linder:2009jz} with $c_{1, 2}$ being constants. The important feature of the exponential gravity is that it has only one more parameter than the $\\Lambda\\mathrm{CDM}$ model. The constraints from the violation of the equivalence principle~\\cite{Tsujikawa:2009ku} and cosmological observations~\\cite{Ali:2010zx} on the exponential gravity have been examined. The exponential gravity in the framework of $f(R)$ gravity has been extended to a gravitational theory in terms of the torsion scalar~\\cite{Linder:2010py} (for a related work on torsion gravity, see~\\cite{Bengochea:2008gz}). We note that the cosmological dynamics in the gravitational theory consisting only of the exponential term without the Einstein-Hilbert one has also been studied in Ref.~\\cite{Abdelwahab:2007jp}.  In this paper, we explicitly investigate the cosmological evolution in the exponential gravity model given by Cognola {\\it et al.}~\\cite{Cognola:2007zu} and Linder~\\cite{Linder:2009jz} in more detail by using the analysis method in Ref.~\\cite{Hu:2007nk}. We also check the above six viability conditions for the model. In particular, we demonstrate that after the matter-dominated stage, the current accelerated expansion of the universe and the crossing of the phantom divide from the phantom phase to the non-phantom (quintessence) one can be realized. It is interesting to note that the crossing of the phantom divide is implied by the cosmological observational data~\\cite{observational status}, while the exponential gravity is a ghost free theory. In addition, we illustrate that the cosmological horizon entropy globally increases with time. We use units of $k_\\mathrm{B} = c = \\hbar = 1$ and denote the gravitational constant $8 \\pi G$ by ${\\kappa}^2 \\equiv 8\\pi/{M_{\\mathrm{Pl}}}^2$ with the Planck mass of $M_{\\mathrm{Pl}} = G^{-1/2} = 1.2 \\times 10^{19}$\\,\\,GeV.  The paper is organized as follows. In Sec.\\ II, we review the model of the exponential gravity in Refs.~\\cite{Cognola:2007zu, Linder:2009jz} and summarize its viability conditions. In Sec.\\ III, we explore the cosmological evolution of the model. We examine the horizon entropy in Sec.\\ IV. Finally, conclusions are given in Sec.\\ V. ",
        "conclusions": "In the present paper, we have studied the cosmological evolution in the exponential gravity. We have summarized various viability conditions and explicitly illustrated that the late-time cosmic acceleration can be realized after the matter-dominated stage. We have also shown that the crossing of the phantom divide from the phantom phase to the non-phantom one can occur and the cosmological horizon entropy globally increases with time. Phenomenologically, at least in the light of the background cosmological evolution, the exponential gravity can be regarded as one of the most promising viable modified gravitational theories because (a) it satisfies all conditions for the viability; (b) in substance it has only one model parameter; and (c) both the current cosmic acceleration following the matter-dominated stage and the crossing of the phantom divide can be realized."
    },
    "1005/1005.5678_arXiv.txt": {
        "abstract": "{We use deep observations taken with the Photodetector Array Camera and Spectrometer (PACS), on board the \\textit{Herschel} satellite as part of the PACS evolutionary probe (PEP) guaranteed project along with submm ground-based observations to measure the dust mass of a sample of high-z submillimeter galaxies (SMGs). We investigate their dust content relative to their stellar and gas masses, and compare them with local star-forming galaxies. High-z SMGs are dust rich, i.e. they have higher dust-to-stellar mass ratios  compared to local spiral galaxies (by a factor of 30) and also compared to local ultraluminous infrared galaxies (ULIRGs, by a factor of 6). This  indicates that the large masses of gas typically hosted in SMGs have already been highly enriched with metals and dust. Indeed, for those SMGs whose gas mass is measured, we infer  dust-to-gas ratios similar or higher than local spirals and ULIRGs. However, similarly to other strongly star-forming galaxies in the local Universe and at high-z, SMGs are characterized by gas metalicities {\\it lower} (by a factor of a few) than local spirals, as inferred from their optical nebular lines, which are generally ascribed to infall of metal-poor gas. This is in contrast with the large dust content inferred from the far-IR and submm data. In short, the metalicity inferred from the dust mass is much higher (by more than an order of magnitude) than that inferred from the optical nebular lines. We discuss the possible explanations of this discrepancy and the possible implications for the investigation of the metalicity evolution at high-z.} ",
        "introduction": "Understanding the evolution of the dust properties and the dust content of galaxies through the cosmic epochs is  crucial for constraining galaxy evolutionary scenarios. The amount of dust has important implications for the star formation (SF). Dust allows the formation of low-mass stars in low-metalicity environments while it inhibits the formation of massive stars, hence affects the IMF \\citep[e.g.][]{omukai05}. Dust greatly enhances the formation of many molecules whose transitions provide the main cooling mechanism for molecular clouds to form stars. Hence  dust is one of the pre-requisites for enhanced SF. Dust also affects  the detectability of high-z galaxies:  dust extinction reduces the detectability of high-z galaxies in the rest-frame UV-optical bands, while dust thermal emission allows the detection of high-z galaxies at mm/far-IR wavelengths. The dust content can also be used as a proxy of the metalicity of the ISM, because refractory elements in the ISM are mostly depleted into dust grains.  The ESA \\textit{Herschel} Space Observatory \\citep{pilbratt10} offers the possibility to investigate the dust emission in high-z star-forming galaxies and in particular to measure their dust content by modeling the far-IR to submm spectral energy distribution (SED). We here present a first investigation of dust mass evolution  by focusing on a sample of high-z submillimeter galaxies (SMGs). They are very luminous ($\\sim 10^{13} L_\\odot$) \\citep[e.g.,][]{hughes98,pope06}, high redshift ($z \\sim$ 1$-$3.5) \\citep{chapman05,pope06}, gas rich and compact, massive galaxies \\citep{greve05,tacconi08}. Their star formation rates (SFR) are exceptionally high ($\\sim 10^3 M_\\odot yr^{-1}$) and are thought to contribute significantly to the cosmic SF at $z \\sim$ 2$-$3 \\citep{chapman05}. Thanks to their high luminosities they provide an excellent laboratory to investigate the evolution of dust in high-z galaxies. To measure their dust content we use data taken with PACS \\citep{poglitsch10} as part of the PACS Evolutionary Probe guaranteed time key program combined with ground-based mm and submm data. By comparing the inferred dust masses with other properties of the SMGs, and with local galaxies, we infer  important information on the physics of these high-z systems. We adopt the $\\Lambda$-CDM concordance cosmological model (H$_0$ = 70 km/s/Mpc, $\\Omega_{\\small M} $ = 0.3 and $\\Omega_{\\Lambda}$ = 0.7). ",
        "conclusions": "Figure~\\ref{fig:mdustmstar} shows the stellar mass vs the dust mass for high-z SMGs (black circles), local ULIRGs (blue squares) and local spirals (red triangles).  The dust masses estimated for local spirals nicely agree with those obtained through the very detailed analysis of \\cite{draine07} of the same sample. Most of them also agree fairly well with the evolutionary pattern for spiral galaxies predicted by the model of \\cite{calura08} (cyan solid line). The dust masses were compared with those obtained by the more common assumption of a single temperature dust distribution. Assuming an emissivity index $\\beta = 2$ (for consistency with GRASIL) and a single, absolute average emissivity of dust at 125$\\mu$m of 2.64 \\citep{dunne03} or 1.87m$^2$ kg$^{-1}$ \\citep{lidraine01}, dust masses are on average 0.3 or 0.15 dex, respectively, lower than those inferred through GRASIL \\citep[100\\mics flux not used in the fitting as in][]{magnelli10}. Independently of the adopted method, there are clear, systematic differences in terms of dust content for the different samples.  The orange dashed and dot-dashed lines in Fig.~\\ref{fig:mdustmstar} show the expected evolution of proto-ellipticals \\citep{calura08} with different mass. These models predict  a larger dust content during the active phase of elliptical formation compared to normal spirals. Nonetheless, they can hardly reproduce the large $\\rm M_{dust}/M_{star}$ observed in most SMGs. The dust-to-star excess in SMGs emerges more clearly in Fig.~\\ref{fig:histo}a, which shows the distribution of $\\rm M_{dust}/M_{star}$  for the different samples. Local ULIRGs have a larger dust content than spirals for a given stellar mass, with  $\\rm M_{dust}/M_{star}$  a factor of $\\sim$4 higher on average. This result is in contrast with \\cite{clements10}, who claimed that the dust content in ULIRGs can be simply explained with the combination of two spirals (the different result  probably arises because Clements et al. use only K-band photometry as a proxy of the stellar mass, while we perform the full SED fitting). Submillimeter galaxies are much more extreme: their $\\rm M_{dust}/M_{star}$  is on average higher by a factor of $\\sim$30 than in local spirals, and by a factor of $\\sim$6  than in local ULIRGs. This large dust content is difficult to account for with the galaxy evolutionary model of \\cite{calura08}.   High-z SMGs are known to be very gas rich, with gas fractions approaching 50\\% \\citep{tacconi08}. The high $\\rm M_{dust}/M_{star}$ values suggest that the large gas masses hosted by SMGs have already been highly enriched with metals and dust; i.e. the excess of dust mass in SMGs is not a consequence of SF episodes following recent inflows of metal-poor gas accreted during mergers \\citep[in contrast with the expectation of some models, ][]{montuori10}. This can be tested by measuring the dust-to-gas ratio for those SMGs with CO millimetric data. By assuming the same CO-to-H$_2$ conversion factor $\\rm \\alpha = 4.3~M_\\odot (K\\ km\\ s^{-1}\\ pc^2)^{-1}$ for all samples (typically applied to normal spirals) the resulting distribution of $\\rm M_{dust}/M_{H_2}$ is shown in Fig.~\\ref{fig:histo}b, similar for the three samples. If one adopts for ULIRGs and SMGs the conversion factor $\\rm \\alpha \\sim 1~M_\\odot (K\\ km\\ s^{-1}\\ pc^2)^{-1}$ \\citep[thought to be more appropriate for these classes of objects, ][]{solomon05,tacconi08}, then ULIRGs and SMGs have $\\rm M_{dust}/M_{H_2}$ significantly larger than local spirals, as indicated by the arrows. It is not possible to infer the total gas mass for SMGs by including the fraction of HI, because this is not observed; but similarly to local ULIRGs, the bulk of the gas is probably in the molecular phase \\citep{sanders96} in contrast to spirals, which may have a considerable fraction of gas in the atomic phase. As a consequence, by including $\\rm M_{HI}$ to obtain the dust-to-total gas ratio, the resulting $\\rm M_{dust}/M_{gas}$ distribution of ULIRGs and SMGs would be further skewed towards higher values relative to spirals. A dust-to-gas ratio higher than in spirals has also been observed in the central (dense) region of local luminous infrared galaxies (LIRGs) \\citep{wilson08}. It can  possibly be ascribed to enhanced dust production by SNe and by the first generation of AGB stars in young stellar systems \\citep[e.g.][]{michalowski10}, and/or to higher accretion of metals on dust grains in these dense environments. In either cases,  the finding that the dust-to-gas ratio in SMGs is {\\it not} lower than in local spirals supports the interpretation of their ``dust richness'' (relative to the stellar mass) in terms of ``gas richness'', possibly along with a higher $\\rm M_{dust}/M_{H_2}$. Summarizing, the high dust content observed in SMGs mirrors their high gas content, the latter with a dust-to-gas ratio similar or higher than normal spiral galaxies.   \\begin{figure}[!t] \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{14748fg2}}  \\caption{\\textit{Panels a), b), d)}: $\\rm M_{dust}/M_{star}$, $\\rm M_{dust}/M_{H_2}$ and $\\rm M_{dust}/M_{Z,gas}$ distributions, respectively, for high-z SMGs (black hatched),  local spirals (SINGS, red open) and local ULIRGs (blue shaded). $\\rm M_{H_2}$ is computed adopting $\\rm \\alpha = 4.3$. Adopting $\\rm \\alpha = 1$ moves the black and blue histograms by the amount indicated by the arrows. \\textit{Panel c)}: metal abundances for the sample of high-z SMGs of \\cite{swinbank04} (black hatched histogram) and that of local ULIRGs of \\cite{rupke08} (blue shaded histogram). Submillimeter galaxies  in common with our sample are marked by arrows. The red solid and green dashed lines show the metalicity distribution for local spirals (SDSS) and $z \\sim 2.2$ star-forming galaxies \\citep{mannucci10}, respectively, with stellar masses similar to those of our SMGs. } \\label{fig:histo} \\end{figure}    The latter result is  in contrast with the low gas metalicity observed in SMGs though. Indeed, based on their optical nebular lines SMGs have metalicities similar or lower than other strongly star-forming systems, such as ULIRGs in the local Universe \\citep{rupke08}, as well as Lyman break galaxies (LBGs), BzKs \\citep[see][]{daddi04} and ULIRGs at high redshift \\citep{erb06,mannucci10,caputi08}, which in turn are lower than local, moderately star-forming spirals with the same stellar mass. This is illustrated in Fig.~\\ref{fig:histo}c, where the metalicity distribution of different classes of objects was measured with the same method (based on optical nebular lines) and adopting the same metalicity scale as discussed in the Appendix, and within similar stellar mass interval ($\\rm M_{\\star}\\sim 10^{11}M_{\\odot}$). For the SMGs we exploited the sample with rest-frame optical spectroscopy given in \\cite{swinbank04}, which unfortunately overlaps with our sample only for three objects (marked with black arrows). That galaxies with enhanced star formation are characterized by lower metalicity has been noted by various authors \\citep{kewley08, rupke08, caputi08}. Recently \\cite{mannucci10} showed that the anticorrelation between metalicity and SFR actually accounts for most of the evolution of the mass-metalicity relation at high redshift: high-z galaxies (at least up to \\textit{z}$\\sim$2.5) appear more metal-poor simply because characterized by higher SFRs as a consequence of selection effects. This anticorrelation between SFR and metalicity is ascribed to inflow of metal-poor gas that boosts SF in galaxies. However, the inferred ``low'' gas metalicities in SMGs are in striking contrast with the  high dust content inferred from their far-IR emission.   In Fig.~\\ref{fig:metal} we further investigate this discrepancy by comparing the dust-to-gas ratio with the gas metalicity. Red triangles are all local galaxies for which the complete required information is available; most of the upper limits are due to the lack of detection or of data for the molecular gas. The dotted line shows the trend expected by assuming that the dust content scales linearly with metalicity following the relation \\citep{draine07} \\begin{equation} \\rm M_{dust}/M_{gas}\\approx 0.01\\times(O/H)/(O/H)_{MW}. \\label{eq:draine} \\end{equation} Actually, \\cite{draine07} and \\cite{hunt05} have shown (on a wider sample) that the relation traced by real galaxies is much steeper than expected by Eq.~\\ref{eq:draine}, as also hinted by the (much smaller) subset of local spirals/dwarfs shown in Fig.~\\ref{fig:metal}. Local ULIRGs (blue symbols) tend to have higher dust-to-gas ratios for a given metalicity, the deviation depending on whether the assumed CO-to-H$_2$ conversion factor is $\\alpha=1$ (open squares) or $\\alpha=4.3$ (crosses). In the case of SMGs (black symbols) there are only two objects for which the complete required information is available (whose gas metalicity is actually an upper limit),  shown both in the case of  $\\alpha=1$ (solid squares) and $\\alpha=4.3$ (solid circles). However, it is clear that their dust content is far higher than expected from their metalicity.  The HI mass is not included for ULIRGs and SMGs, but the effect is very small \\citep[0.1 dex, ][]{sanders96}.  The orange and cyan tracks in Fig.~\\ref{fig:metal} show the expected evolution of the dust content and metalicity as predicted by \\cite{calura08} for ellipticals and spirals, where it is evident that the large discrepancy observed in SMGs cannot be accounted for by models. In short, in SMGs the mass of metals inferred from the dust content is much larger than the mass of metals inferred from  the gas phase.  To further quantify this discrepancy, in Fig.~\\ref{fig:histo}d we show the distribution of $\\rm M_{dust}/M_{Z,gas}$, where $\\rm M_{Z,gas}$ is the mass of metals in the gas phase (i.e. $\\rm Z\\times M_{gas}$). We  assumed for all objects a conversion factor $\\alpha=4.3$, the arrows showing the effect of using $\\alpha=1$ for ULIRGs and SMGs. Clearly in SMGs $\\rm M_{dust}/M_{Z,gas}$ is over one order of magnitude higher than in spirals and well above unity. Local ULIRGs share a similar problem, but not as extreme as SMGs. When computing $\\rm M_{Z,gas}$, we assume the same metal abundance for the neutral  and  ionized  gas phases. This can be false, as the former may be characterized by  lower \\citep[by a factor of $\\sim$7, ][]{leboutillier09} abundances than the latter. However, this effect does not change our results. Indeed, dwarf low-metalicity galaxies (see the Appendix), where most of the gas is in the atomic phase, would move towards higher values of  $\\rm M_{dust}/M_{Z,gas}$ by 0.85 dex, while normal spiral galaxies, whose atomic gas fraction is around 50\\%, would only move by 0.24 dex. The ratio  $\\rm M_{dust}/M_{Z,gas}$ in spiral galaxies and dwarfs would still be lower than in  SMGs, where the bulk of the gas is expected to be in the molecular phase (as discussed above).  There are some possible scenarios that could explain this puzzling result. One possibility is that dust masses are erroneously estimated because the dust properties in SMGs are very different from those assumed in our models. The average dust emissivity may be higher. However, this difference would have to be quite dramatic, because the dust masses are off by more than an order of magnitude with respect to those expected by the metalicities. In ULIRGs (extreme starbursts) the emissivity inferred by \\cite{clements10} is similar to that assumed here.  Another possibility is that as mentioned above the high density in SMGs favors the growth of dust out of metals in the ISM, which also decreases the effective metalicity in the gas phase. However, the standard dust depletion factors in the diffuse ISM are already high, with most metals already locked into dust. Hence, for a given metalicity the dust mass cannot grow by more than a factor of about two through this mechanism.  Finally, the disagreement between the dust content and the gas metalicity can be caused by dust obscuration. Likely, the bulk of the gas in SMGs is metal rich. The associated dust richness and the compact configuration make most of the ISM optically thick at visual wavelengths. As a consequence, the optical nebular lines used to infer the gas metalicity only probe the outer parts of the star-forming regions, which are probably more metal-poor. Paradoxically, the higher metalicity of these dense systems and their associated high dust content may produce the apparent effect of  metal-poorness when observed at optical wavelengths. This scenario is supported by the findings of \\cite{santini09}, who showed that luminous IR galaxies have higher mid-IR-based SFR estimates compared to those obtained by correcting the optical-UV light.  If confirmed, the latter two scenarios would not support the current interpretation of the low metalicity observed in ULIRGs and SMGs in terms of inflow of metal-poor gas boosting the star formation \\citep[e.g.][]{montuori10}. More generally, if the discrepancy between gas and dust metal mass also applies  to other populations of high-z starburst galaxies (e.g. BzKs and LBGs) this would prompt  a general revision of the interpretation of the apparently low metalicities observed in high-z systems.      \\begin{figure}[!t] \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{14748fg3}} \\caption{ Dust-to-gas ratio versus metalicity. The black and blue symbols show high-z SMGs and local ULIRGs, respectively, by assuming  different  conversion factors  as indicated by the legend. Red triangles refer to local spirals. The green dotted line shows the trend expected by Eq.~\\ref{eq:draine}. The solid cyan and dashed (dot-dashed) orange lines show the evolutionary tracks predicted by the model of  \\cite{calura08} for spirals and proto-ellipticals with mass of $\\rm 10^{11}$ ($10^{12}$) $\\rm M_{\\odot}$.} \\label{fig:metal} \\end{figure}"
    },
    "1005/1005.0788_arXiv.txt": {
        "abstract": "We construct cosmological models consisting of large numbers of identical, regularly spaced masses.  These models do not rely on any averaging procedures, or on the existence of a global Friedmann-Robertson-Walker (FRW) background.  They are solutions of Einstein's equations up to higher order corrections in a perturbative expansion, and have large-scale dynamics that are well modelled by the Friedmann equation. We find that the existence of arbitrarily large density contrasts does not change either the magnitude or scale of the background expansion, at least when masses are regularly arranged, and up to the prescribed level of accuracy. We also find that while the local space-time geometry inside each cell can be described as linearly perturbed FRW, one could argue that a more natural description is that of perturbed Minkowski space (in which case the scalar perturbations are simply Newtonian potentials). We expect these models to be of use for understanding and testing ideas about averaging in cosmology, as well as clarifying the relationship between global cosmological dynamics and the static space-times associated with isolated masses. ",
        "introduction": "The Friedmann-Robertson-Walker (FRW) model is ubiquitous in modern cosmology, and is widely believed to be a good model for a universe with a matter content that is approximately uniformly distributed over large scales.  FRW has had great success in fitting a variety of cosmological observations, but is not without its problems. In particular, it has the apparent defect of requiring large amounts of dark matter and dark energy in order to be compatible with observations. More fundamentally, it relies on the implicit assumption that it is permissible to use non-local `average' energy densities in Einstein's equations, which are a set of local field equations. The difficulties involved with this are exacerbated by the diffeomorphism covariance of the theory, which means that there is in general no preferred set of space-like hypersurfaces with which to perform an average over at all.  What is more, even if a suitable and unique averaging scheme is found, it seems we will still be left with a back-reaction effect due to the non-commutativity of averaging and evolution under Einstein's equations \\cite{ellis,buchert}.  These issues require further study in order to be fully understood.  One way to make progress in this area is to construct alternative models that are approximately homogeneous and isotropic on the largest scales, but that do not involve any averaging procedures.  We will survey some of the relevant literature on progress toward this goal in the section that follows, and then proceed to present our own approach to constructing a cosmological model that is composed of discrete masses, rather than a continuous fluid.  This model will appear homogeneous and isotropic when coarse-grained over the largest scales, but will also be able to accommodate arbitrarily large local density contrasts.  The validity of this model will be based solely on Newtonian and post-Newtonian perturbative expansions about Minkowski space, which are generally considered to be able to easily model large density gradients without breaking down (for a recent discussion of the potential difficulties involved in modelling the corresponding situation with a perfect fluid filled, perturbed FRW cosmology see, e.g., \\cite{ras}).  The model we construct is one in which isolated masses are arranged in a regular array, which by application of the Israel junction conditions is a solution of Einstein's equations up to a specified level of accuracy.  Contrary to some previous results in the literature, we show that discretisation of the matter content does {\\it not} affect the global expansion of the space-time, which proceeds in just the same way as a perfect fluid FRW solution with the same energy density and spatial curvature, up to the required accuracy.  This result says nothing about what happens in the more realistic situation of irregularly arranged masses, but does show that, in at least some situations, perturbed FRW is a good description of a universe with (arbitrarily) large density contrasts.  Establishing the existence of such situations is of interest not only for limiting the possible consequences of inhomogeneity on the large-scale expansion of the Universe, but also for testing the viability of proposed methods of accounting for more general inhomogeneity in the Universe:  If they predict modifications to the large-scale expansion in configurations where it is known that none should occur, then one may choose to question their viability.  In Section \\ref{2b} we survey some of the literature on inhomogeneous cosmological models, concentrating in particular on studies that have similarities to the approach used here, and highlighting predictions of deviations from the usual Friedmann expansion.  In Section \\ref{3} we present the type of model we will be considering, and the perturbative expansion we deploy is outlined in Section \\ref{pertsec}.  Section \\ref{statsec} treats the geometry inside each of the primitive cells of our structure as being perturbed Minkowski space, and finds the corresponding cosmological evolution, after the Israel junction conditions have been applied.  A similar analysis is then performed in Section \\ref{timesec}, this time around a time dependent background.  In Section \\ref{7} we relate the time independent and time dependent approaches used in the preceding two sections, and in Section \\ref{8} we show that the model we have constructed has identical large-scale expansion to a perfect fluid filled FRW cosmology with the same energy density and large-scale spatial curvature.  In Section \\ref{obs} we consider the problem of determining cosmological observables, such as redshifts and luminosity distances, and in Section \\ref{disc} we conclude. ",
        "conclusions": "\\label{disc}  In summary, we have considered n-body cosmological models that do not require any averaging procedures.  These models have been constructed using a lattice made from a large number of symmetric cells, each containing identical central masses.  By applying appropriate junction conditions between cells we then find a set of evolution equations that specify the motion of the cell boundaries, and hence the global evolution of the space-time. In all of the cases we have considered we find that the cell boundaries must be in free fall, and decelerating in their expansion away from the cell centres.  We find that the space-time geometry inside each cell can be described as fluctuations around either static or time-dependent FRW backgrounds.  In order to apply the FRW description one must include potentials of the form $\\Phi \\sim r^2$, whose gradients must be of the same order of magnitude as the Newtonian potentials that are also present.  Treating the space-time inside each cell as fluctuations around Minkowski space, however, requires Newtonian potentials only.  One may then consider perturbed Minkowski space to be a more natural local description of the space-time.  We have compared the resulting large-scale dynamics of the cosmological model under consideration with those of a perfect fluid FRW cosmology.  We find that for spatially flat universes the expansion rate for a given energy density has just the expected FRW value. Furthermore, we have compared the scale of expansion of non-flat models with spatially curved perfect fluid FRW universes and found that they are also the same, up to the required accuracy.  We emphasize that this result does {\\it not} follow directly from the fact that the space-time geometry inside each cell can be written as perturbed FRW: It could, in principle, have been the case that the scale of curvature of the global lattice was different from the scale inferred from the relevant FRW background used inside each cell (in fact, from a number of the previous studies discussed in Section \\ref{2b}, it would seem that one may have expected such a result). Instead, what we find here is that discretisation of the matter content of the Universe does {\\it not} have to affect the global background rate of expansion.  These results are a consequence of explicitly solving boundary conditions between different regions that can, individually, be described as perturbations around highly symmetric backgrounds.  That such a description is possible within each cell does not, however, automatically mean that one can treat the entire Universe as fluctuations around a single universal FRW background.  Such a result has to be shown to be true by reconstructing the global geometry from the local geometry that is appropriate within each of these regions.  This is what we have done here, finding that for the simple case of regularly spaced masses a single global FRW background with small perturbations around it is a perfectly adequately description of the entire space-time.  That a global FRW background is a valid description in the present case, however, does not automatically mean that this will also be true for irregular arrangements of massive objects.  Such a result remains to be shown.  As already discussed, the model we have considered here can be interpreted solely in terms of perturbations about Minkowski spaces (albeit a different Minkowski space within each cell).  That this is possible gives one greater reason to expect our perturbative expansion to be valid in the regime of non-linear density contrasts than is the case in globally FRW perfect fluid cosmology.  As discussed in \\cite{ras}, there are terms at higher orders in the latter case that one may expect to blow up when $\\delta \\rho/\\rho$ becomes large.  With the former case of Newtonian and post-Newtonian fluctuations about Minkowski space, however, the situation is different.  The terms that could potentially blow up are absent, and hence we have greater confidence in the applicability of the post-Newtonian description of these systems when density contrasts are large.  To determine whether or not the results we have found here carry over to more realistic models of the Universe, where masses are irregularly spaced, and dispersed matter is also present, will require further study, and more refined models.  However, we do expect our results to be of use for constraining the possible effects that structure formation could have on the background expansion of the Universe, as well as for testing the viability of mechanisms that have been constructed in order to correct for back-reaction and averaging:  If corrections are predicted in situations where it is known that none occur, then the frame-work within which they have been identified should be questioned.    \\newpage  \\vspace{10pt}  \\begin{flushleft} {\\bf Acknowledgements} \\end{flushleft}  I am grateful for valuable discussions with Pedro Ferreira, Chris Clarkson, George Ellis and Thomas Sotiriou.  I also wish to acknowledge the support of Jesus College Oxford, the BIPAC, and CERN."
    },
    "1005/1005.2572_arXiv.txt": {
        "abstract": "s{Polarized foregrounds are going to be a serious challenge for detecting CMB cosmological B-modes. Both diffuse Galactic emission and extragalactic sources contribute significantly to the power spectrum on large angular scales. At low frequencies, Galactic synchrotron emission will dominate with fractional polarization $\\sim 20-40\\%$ at high latitudes while radio sources can contribute significantly even on large ($\\sim 1^{\\circ}$) angular scales. Nevertheless, simulations suggest that a detection at the level of $r=0.001$ might be achievable if the foregrounds are not too complex. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.2434_arXiv.txt": {
        "abstract": "Time-resolved spectra throughout the orbit of EF Eri during its low accretion state were obtained with the Solar Blind Channel on the Advanced Camera for Surveys onboard the {\\it Hubble Space Telescope}. The overall spectral distribution exhibits peaks at 1500 and 1700\\AA, while the UV light curves display a quasi-sinusoidal modulation over the binary orbit. Models of white dwarfs with a hot spot and cyclotron emission were attempted to fit the spectral variations throughout the orbit. A non-magnetic white dwarf with a temperature of $\\sim$10,000K and a hot spot with central temperature of 15,000K generally matches the broad absorptions at 1400 and 1600\\AA\\ with those expected for the quasimolecular H features H$_{2}$ and H$_{2}^{+}$. However, the flux in the core of the Ly$\\alpha$ absorption does not go to zero, implying an additional component, and the flux variations throughout the orbit are not well matched at long wavelengths. Alternatively, a 9500K white dwarf with a 100 MG cyclotron component can fit the lowest (phase 0.0) fluxes, but the highest fluxes (phase 0.5) require an additional source of magnetic field or temperature. The 100 MG field required for the UV fit is much higher than that which fits the optical/IR wavelengths, which would support previous suggestions of a complex field structure in polars. ",
        "introduction": "EF Eri is a well-studied cataclysmic variable that contains a highly magnetic white dwarf. It was identified in 1979 as the fourth known Polar (Griffiths et al. 1979; Hiltner et al. 1979; Tapia 1979) with an 81 min orbital period. During the next 20 years, it remained in a bright high state near 14th mag, with highly modulated optical, IR and X-ray light curves due to a high mass transfer rate and X-ray and cyclotron emission from the accretion region (Bailey et al. 1982). Models of the cyclotron humps determined magnetic field strengths of 16.5 and 21 MG (Ferrario, Bailey \\& Wickramasinghe 1996). However, in 1997, EF Eri entered a low state near 18th mag (Wheatley \\& Ramsay 1998) which lasted for 9 years, with only short excursions to a high state since 2006. Low states are advantageous to study the underlying white dwarf and secondary star without the effects of the accretion flow. In the low state, the optical spectrum of EF Eri shows broad Balmer absorption lines with Zeeman splitting indicating a 14 MG field (Wheatley \\& Ramsay 1998). Beuermann et al. (2000) modeled the optical spectrum with a 9500K white dwarf and a 15,000K hot spot over 6\\% of its surface. As there were no features from a secondary star, they concluded that the spectral type of the companion must be later than M9. {\\it Spitzer} mid-IR emission suggested the companion was an L or T dwarf (Howell et al. 2006b). Howell et al. (2006a) used the H$\\alpha$ emission from this substellar secondary to produce a radial velocity curve, which provided a spectroscopic ephemeris as well as a mass estimate for the secondary of 0.055 M$_{\\odot}$ for a white dwarf mass of 0.6 M$_{\\odot}$. Thorstensen (2003) attempted parallax measurements but the faintness of the source and the assumptions used yielded a small range of distances centered near 113 pc or a larger range near 163 pc.  While the accretion rate in low states is reduced by several orders of magnitude from the high state, and there is no evidence of a stream of mass transfer during the low state of EF Eri, UV studies have shown that hot spots with temperatures of 30,000-70,000K are likely present on the white dwarfs of Polars during low states. {\\it IUE} observations of AM\\,Her in the low state revealed an orbital-phase dependent flux modulation, which Heise \\& Verbunt (1988) interpreted as the signature of moderate heating of a large fraction of the white dwarf surface. G\\\"ansicke et al. (1995) developed this hypothesis further, showing that a large pole cap covering $\\sim10\\%$ of the white dwarf surface is present both during the high state and the low state, and that, at least in the case of AM\\,Her, the luminosity of this heated pole cap is comparable to that of the hard X-ray and cyclotron emission from the post-shock region, solving the so-called ``soft X-ray puzzle'', i.e. the soft X-ray excess with respect to the simple reprocessing model of Lamb \\& Masters (1979) and King \\& Lasota (1979). Evidence for large, ``warm'' pole-caps was found in many other polars (e.g. Stockman et al. 1994; G\\\"ansicke et al. 2000, Schwope et al. 2002, Araujo-Betancor et al 2005). G\\\"ansicke et al (1998) developed a 3-dimensional model of a white dwarf with a heated pole cap to analyze high-time resolution \\textit{HST}/GHRS observations of AM\\,Her in a high state, and \\textit{FUSE} and \\textit{HST}/STIS observations of AM\\,Her in a low state (G\\\"ansicke et al 2006), leading to more precise constraints on the size, temperature, and location of the heated pole cap. Especially for the low-state data, where no emission from the accretion stream was present, this model provided an excellent fit to both the observed modulation of the ultraviolet flux, as well as to the changes in the Lyman line profiles.  {\\it GALEX} photometric observations of EF Eri during its low state (Szkody et al. 2006; 2008) in two broad band filters centered near 1550\\AA\\ and 2300\\AA\\ were modeled with a 9500K white dwarf and a spot of radius 5.5$^\\circ$ with a temperature of 24,000K. Schwope et al. (2007) analyzed archival {\\it XMM-Newton} observations to detect an X-ray flux of 6$\\times$10$^{-14}$ erg cm$^{-2}$ s$^{-1}$ during the low state, corresponding to an x-ray luminosity of 2$\\times$10$^{29}$ erg s$^{-1}$. While the specific accretion rate would then be only 0.01 g cm$^{-2}$ s$^{-1}$, implying no accretion shock, they suggested the X-rays come from some residual accretion at the poles. Li et al. (1994, 1995) have produced models for high field white dwarfs where the white dwarf can capture all the stellar wind from the secondary and funnel it to the poles.  Further support for some accretion during the low state comes from the presence of cyclotron humps in the IR (Harrison et al. 2004; Campbell et al. 2008). Campbell et al. (2008) explored whether high magnetic field cyclotron emission could be producing the {\\it GALEX} UV variations as well as the IR. The high field (240 MG) system AR UMa does show UV variability that is likely due to cyclotron emission (G\\\"ansicke et al. 2001) and QS Tel showed humps at 2200\\AA\\ and 2430\\AA\\ that could be interpreted as cyclotron humps in a 65 MG field (Rosen et al. 2001). The models of Campbell et al. (2008) showed that the UV variations in EF Eri could be cyclotron if the field was 115 MG, while the IR variations require a field of 13 MG. Some confirmation of the high field required for the UV is evident from the work of Beuermann et al. (2007), who used Zeeman tomography and determined that EF Eri has a complex field with some regions as high as 100 MG.  In order to better understand the cause of the UV light and variability, we obtained time-resolved spectroscopy with the {\\it Hubble Space Telescope (HST)}. These data and our modeling efforts with a heated polecap and cyclotron emission are described below. ",
        "conclusions": "Our low resolution {\\it HST} spectra have revealed the first UV spectrum of EF Eri during its low accretion state and the corresponding changes in the spectra throughout its orbit. Two interpretations of the lowest orbital UV fluxes at phase 0.0 are possible. If most of the flux is from a white dwarf at a distance of 120 pc, the UV spectrum can be matched with a temperature of $\\sim$10,000K and a hot spot with central temperature of 15,000K (similar values found by Beuermann et al. (2000) for the optical region). The broad absorptions at 1400 and 1600\\AA\\ match those expected for the quasimolecular H features H$_{2}$ and H$_{2}^{+}$ but the flux in the core of Ly$\\alpha$ does not go to zero, implying an additional component. On the other hand, the UV spectrum can also be adequately modeled with a 9500K white dwarf at 120 pc that contributes longward of 1600\\AA\\ in combination with a high (100MG) cyclotron component dominating at the shortest wavelengths.  Both models have difficulty reproducing the observed orbital flux variations, especially for wavelengths longer than 1400\\AA. A normal white dwarf with a hot spot does not produce enough amplitude near 1500\\AA\\, but produces too much flux at 1700\\AA\\ to account for the phase 0.5 modulation. This type of model has problems producing the amplitudes of the optical flux variations as well. The cyclotron models also require some additional component (either in temperature or in additional field strengths) to account for the phase 0.5 UV fluxes longward of 1400\\AA. The failure of the simple Constant Lambda model to fit all phases points to the need for more complex multi-temperature, multi-lambda models. While our model excludes a simple UV emitting pole, a more complex accretion region is very possible. The Zeeman tomography work by Beuermann et al (2007) on EF Eri, BL Hyi and CP Tuc shows that the field structures are not simple dipoles and there is a range  of field strengths in all three of these polar systems. For EF Eri, the multipole fields range from  4 to 111 MG.  This range could accomodate the fields of 13-21 MG that optimally fit the optical and IR spectral regions (Ferrario et al. 1996; Campbell et al. 2008) as well as the UV field strength estimated from ${\\it GALEX}$ (Szkody et al. 2008) and our HST data. An additional uncertainty is related to the distance. Ultimately, the choice of model would be best served by UV spectropolarimetry to define the cyclotron component correctly, but a good parallax and UV spectra can pin down the UV luminosity from a high field accreting white dwarf. The theoretical profiles for Ly$\\alpha$ and the quasimolecular H features at high fields are not well-known, so observations are currently the best way to determine the influence of the field on the resulting radiation."
    },
    "1005/1005.5508_arXiv.txt": {
        "abstract": "In this paper,  the cosmological  dynamics of Brans-Dicke theory in which there are fermions with a coupling to BD scalar field as well as a self-interaction potential is investigated. The conditions that there exists a solution which is stable and represents a late-time accelerated expansion of the universe are found. The variable mass of fermions can not vanish exactly during the evolution of the universe once it exists initially.  It is shown that the late-time acceleration depends completely on the self-interaction of the fermion field if our investigation is restricted to the theory with positive BD parameter $\\omega$. Provided a negative $\\omega$ is allowed, there will be another two class of stable solutions describing late-time accelerated expansion of the universe. ",
        "introduction": "Strong evidences from the current cosmological observations such as supernovae type Ia(SNeIa)\\cite{A.G. Riess}, cosmic microwave background (CMB)\\cite{D. N. Spergel et al.} and large scale structure(LSS)\\cite{M. Tegmark et al.} converge upon the fact that the universe is spatially flat and there exists exotic component, dubbed dark energy, which drives the speed-up expansion of the universe. Many scenarios have been proposed to explain the acceleration in the framework of general relativity (GR). The preferred and simplest candidate for dark energy is the Einstein's cosmological constant which can fit the observations well. However, it suffers from the so-called fine-tuning problem  and coincidence problem. On the other hand, the observations are not yet able to confirm that dark energy is indeed a constant. Actually,  many dynamical models of dark energy have been studied extensively during the past over ten years, such as quintessence \\cite{Quintessence}, phantom \\cite{Phantom}, quintom \\cite{Quintom} (see also Ref.\\cite{cai} for a detailed review),  tachyon \\cite{tacyon}, Chaplygin gas \\cite{cg}\\footnote{It has been shown that only the generalized version of the Chaplygin gas model is compatible with the observations (see e.g. Ref. \\cite{gcg}).}, etc (see Ref.\\cite{copeland:2006ej} for a comprehensive review on dark energy models).   GR is a well tested theory in solar system, but it is poorly tested on cosmic scales. Therefore, it is interesting to ask whether the cosmic acceleration is caused by a modified theory of gravity. As is pointed out, the dark energy problem may be essentially an issue of quantum gravity \\cite{Witten:2002}. Although a complete  theory of quantum gravity has not been established, some valuable ideas are thought to be the features of the theory of quantum gravity, for example, extra dimensions. Many fundamental theories would induce scalar fields in the usual 4-dimensional space-time, and these scalar fields are in general non-minimal coupled. It is truly remarkable to find that a candidate scalar field of the desired nature is provided by some fundamental theory\\footnote{Note that string theory is a quite special type of scalar-tensor theory of gravity, and unless one introduces non-perturbative effects, the theory has a  runaway problem. }. On the other hand, in recent years, the interest in scalar-tensor theories of gravity as viable alternatives to GR is renewed. In particular, some authors have resorted to scalar-tensor theory in order to explain the present accelerating expansion of the universe \\cite{Bertolami:2000,Brans-Dicke:DE}.  As the simplest example and prototype of scalar-tensor theory of gravity, Brans-Dicke theory (BD) formulates the gravitational phenomena through the interplay between the metric tensor and a scalar field $\\phi$ that controls the intensity of the gravitational constant $G$, while the coupling between $\\phi$ and matter is absent. Although there is no direct support from fundamental theories, it is  hard to deny that it appears to provide a small window through which one can look into phenomenological aspects of fundamental theories.  Recently, the possibility that fermion fields as gravitational sources could be responsible for accelerated periods during the expansion of the universe is considered in the framework of GR\\cite{Ribas:2005vr} as well as in BD theory of gravity\\cite{Samojeden:2010rs}. In the context of cosmology, due to the isotropy and homogeneity of the geometry of the universe, the fermion field configurations should be constant in space. It is shown that the fermion field could behave as dark energy for late-time universe. In this work, we would like to investigate the  dynamics of BD cosmology in which there exists fermions with coupling to BD scalar field as well as a self-interaction potential, and the main focus is looking for the conditions under which there is a solution that is stable and represents a late-time accelerated expansion of the universe.  This paper is organized as follows: In section \\ref{sec2}, we briefly introduce the tetrad formalism of fermions with variable mass in Brans-Dicke theory of gravity. Then, the cosmological dynamics of the model is investigated in detail in section \\ref{sec3}. Finally, in section \\ref{sec:conclusion}  we summarize our results and give some discussions. ",
        "conclusions": "\\label{sec:conclusion} Brans-Dicke theory is a natural alternative and a simple extension of general relativity. In this work, we have studied the cosmological dynamics of Brans-Dicke theory in which there are fermions with a coupling to BD scalar field as well as a self-interaction potential. Because of the isotropy and homogeneity of the geometry of the universe, the fermion field is independent of space and becomes an exclusive function of time.  The conditions that there exist solutions which are stable and describe a late-time accelerated expansion of the universe are found. The variable mass of fermions can not vanish exactly during the evolution of the universe once it obtains initially.  It is shown that the late-time acceleration depends completely on the self-interaction of the fermion field if our investigation is restricted to the theory with positive BD parameter $\\omega$. Provided a negative $\\omega$ is allowed, there will be another two class of stable solutions describing late-time accelerated expansion of the universe which are independent of the self-interaction of fermion field.  The strongest constraint to date on the Brans-Dicke theory of gravity has been put on the solar system scale\\cite{Bertotti:2003b}, however, it should be pointed out that such constraint does not necessarily apply on scales much larger than those of the measurements, and on epochs much different from the present. The solar system experiments only probe scales in gravitational equilibrium, where the background expansion of the universe is negligible, therefore, they would not reveal spatial or time variation of the gravitational constant on larger scales. Through a certain  mechanism, the BD parameter $\\omega$ becomes variable from small scales to large scales.  It is even not excluded that on the largest scales the gravitational constant becomes effectively negative. Some indications for a negative $\\omega$ (or effectively negative gravitational constant) have  been obtained \\cite{Bertolami:2000}. Therefore, it is interesting to investigate such a model presented here in the context of a detailed viable gravitational theory with variable effective $\\omega$, for example, the abnormally weighting energy theory \\cite{AWE}."
    },
    "1005/1005.3118_arXiv.txt": {
        "abstract": "{The \\emph{Herschel} OB young stellar objects survey (HOBYS) has observed the Rosette molecular cloud, providing an unprecedented view of its star formation activity. These new far-infrared data reveal a population of compact young stellar objects whose physical properties we aim to characterise. We compiled a sample of protostars and their spectral energy distributions that covers the near-infrared to submillimetre wavelength range. These were used to constrain key properties in the protostellar evolution, bolometric luminosity, and envelope mass and to build an evolutionary diagram. Several clusters are distinguished including the cloud centre, the embedded clusters in the vicinity of luminous infrared sources, and the interaction region. The analysed protostellar population in Rosette ranges from 0.1 to about 15\\,\\msun\\ with luminosities between 1 and 150\\,\\lsun, which extends the evolutionary diagram from low-mass protostars into the high-mass regime. Some sources lack counterparts at near- to mid-infrared wavelengths, indicating extreme youth. The central cluster and the Phelps~\\&~Lada~7 cluster appear less evolved than the remainder of the analysed protostellar population. For the central cluster, we find indications that about 25\\% of the protostars classified as Class\\,I from near- to mid-infrared data are actually candidate Class\\,0 objects. As a showcase for protostellar evolution, we analysed four protostars of low- to intermediate-mass in a single dense core, and they represent different evolutionary stages from Class\\,0 to Class\\,I. Their mid- to far-infrared spectral slopes flatten towards the Class\\,I stage, and the 160 to 70\\,\\micron\\ flux ratio is greatest for the presumed Class\\,0 source. This shows that the \\emph{Herschel} observations characterise the earliest stages of protostellar evolution in detail.} ",
        "introduction": "The HOBYS \\emph{Herschel} \\citep{pilbratt2010} imaging survey \\citep{motte2010} is a key programme for studying the sites of OB star formation within a distance of 3\\,kpc. The parallel mode scan map observations, carried out with the PACS \\citep{poglitsch2010} and SPIRE \\citep{griffin2010,swinyard2010} instruments, provide maps of massive cloud complexes in the wavelength range of 70 to 500\\,\\micron, and thus a census of the different stages of star formation from prestellar cores to evolved young stellar objects (YSOs). As an excellent example of a massive cloud complex associated with -- and under the influence of -- an OB star cluster (NGC\\,2244), the Rosette molecular cloud was observed for HOBYS during the \\emph{Herschel} Science Demonstration Phase \\citep{motte2010}. The influence of NGC\\,2244 on the cloud complex is discussed in the accompanying paper by \\citet{schneider2010}, while \\citet{difrancesco2010} assesses the clump population. For consistency with \\citet{schneider2010}, the distance of Rosette adopted here is 1.6\\,kpc. Many previous studies have targeted the Rosette complex. Ongoing star formation across the molecular cloud is traced by the presence of several luminous IRAS sources that are associated with massive clumps seen in CO emission \\citep{cox1990,williams1995,schneider1998}. Towards these dense regions, the embedded clusters PL1 to PL7 have been identified by \\citet{phelpslada1997} from near-infrared observations. The list of clusters in Rosette was extended by \\citet{lismith2005}, \\citet{romanzuniga2008}, and \\citet{poulton2008} using near-infrared and \\emph{Spitzer} observations.  In this paper, we make use of the unprecedented spatial resolution and sensitivity of \\emph{Herschel} to establish a sample of compact far-infrared sources that represent protostellar objects and to determine their fundamental properties (luminosities and envelope masses) from their spectral energy distribution (SED). In the early phases of the collapse of a protostellar core and the initial accretion of matter onto a central protostar, the SED of a Class\\,0 source is dominated by thermal emission from cold dust in the envelope \\citep[$\\menv > M_\\star$,][]{andre2000}. The SED of a more evolved Class\\,I source is shifted to mid-infrared wavelengths, indicating a comparatively less massive, hotter envelope ($\\menv<M_\\star$). To investigate the evolutionary stage, the protostellar envelope mass is usually compared to the bolometric luminosity, which serves as a proxy for stellar mass. The \\emph{Herschel} observations for the first time cover the peak of the protostellar SED thus constraining the evolutionary stage of early-to-evolved YSOs. This overcomes the previous difficulties distinguishing Class\\,0 and Class\\,I sources using near- to mid-infrared data. ",
        "conclusions": ""
    },
    "1005/1005.3604_arXiv.txt": {
        "abstract": "We present a multi-wavelength analysis of the merging rich cluster of galaxies Abell 2256. We have observed A2256 at 150 MHz using the Giant Metrewave Radio Telescope and successfully detected the diffuse radio halo and the relic emission over an extent $\\sim1.2$ Mpc$^2$. Using this 150 MHz image and the images made using archival observations from the VLA (1369 MHz) and the  WSRT (350 MHz), we have produced spectral index images of the diffuse radio emission in A2256. These spectral index images show a distribution of flat spectral index (S$\\propto\\nu^\\alpha$, $\\alpha$ in the range -0.7 to -0.9) plasma in the NW of the cluster centre. Regions showing steep spectral indices ($\\alpha$ in the range -1.0 to -2.3) are toward the SE of the cluster centre. These spectral indices indicate synchrotron life times for the relativistic plasmas in the range 0.08 - 0.4 Gyr. We interpret this spectral behaviour as resulting from a merger event along the direction SE to NW within the last 0.5 Gyr or so. A shock may be responsible for the NW relic in A2256 and the Mpc scale radio halo towards the SE is likely to be generated by the turbulence injected by mergers. Furthermore, the diffuse radio emission shows spectral steepening toward lower frequencies. This low frequency spectral steepening is consistent with a combination of spectra from two populations of relativistic electrons created at two epochs (two mergers) within the last $\\sim$0.5 Gyr. Earlier interpretations of the X-ray and the optical data also suggested that there were two mergers in Abell 2256 in the last 0.5 Gyr, consistent with the current findings. Also highlighted in this study is the futility of correlating the average temperatures of thermal gas and the average spectral indices of diffuse radio emission in respective clusters. ",
        "introduction": "The intra-cluster medium (ICM), which pervades the space between the galaxies in galaxy clusters is known to consist of hot thermal gas ($\\sim 10^8$ K), magnetic fields ($\\sim 1\\mu$G) and relativistic particles. The thermal gas emits mainly in X-rays by thermal Bremsstrahlung mechanism. A direct evidence for magnetic fields and relativistic electrons in the ICM are the extended diffuse synchrotron sources associated with the ICM detected in a fraction of clusters (Ferrari et al 2008). These are classified broadly as radio halos (centrally located in clusters and unpolarized ($\\lesssim 5\\%$)) and radio relics (filamentary/ arc-like, located at cluster peripheries and polarized ($\\sim20-30\\%$)) (Ferrari et al 2008). So far, such sources have been detected in merging clusters of galaxies; thus a strong connection between their origin and merger is favoured (Ferrari et al 2008). The primary models predict that shocks and/or turbulence induced during merging events reaccelerate the electrons in the ICM: shocks reaccelerate by Fermi I process (Ensslin et al 1998) and/or adiabatic compression of fossil radio plasma (Ensslin \\& Gopal-Krishna 2001); turbulence reaccelerate via stochastic, Fermi II or MHD waves (Brunetti et al 2004; Cassano \\& Brunetti 2005). The secondary models regard the synchrotron emission to come from electrons generated in hadronic collisions (Dolag \\& Ensslin 2000 and references therein); this model awaits observational evidence and is not discussed here. \\par The shocks and/or turbulence in the ICM which are responsible for accelerating charged particles leave signatures in the spectra of accelerated particles. A shock can leave a trail of accelerated charged particles in its wake; turbulence can lead to patchy distribution of spectral index in the source (Feretti et al 2004 and references therein). Synchrotron spectral index steepens with time as a result of energy losses proportional to E$^2$. In a simplistic approach, the spectral index can be used to estimate the time since the relativistic plasma was accelerated-- the spectral age. If the break frequency of the spectrum can be identified, the spectral age can be estimated; if it cannot be identified, at least a limit on the spectral age can be obtained. Spectral age distribution across the extents of radio halos and relics can then be connected to the geometry of shock passage or the sites of efficient/inefficient turbulent acceleration. Due to the complex nature of the reacceleration by turbulence, this may not be the best estimate of the time since the acceleration took place. Consistency with the merger geometries proposed based on X-ray surface brightness and optical galaxy distributions can be verified. Such studies require good quality multi-frequency maps of radio halos and relics. Radio halos and relics being extended ($\\sim5$ to $30$ arcminutes for clusters in redshift range $0.2 - 0.02$), low surface brightness ($\\sim$ 1 mJy arcmin$^{-2}$ at 1.4 GHz) sources are difficult to image at multiple frequencies with comparable sensitivities. Spectral index maps have been constructed only in a few clusters so far (for example, Coma (Giovannini et al 1993), A665, A2163 (Feretti et al 2004), A3562 (Giacintucci et al 2005), A2744 and A2219 (Orru et al 2007) and A2255 (Pizzo \\& Bruyn 2009)). These have shown features like radial spectral steepening and patchy distribution of spectral indices and have been interpreted as variations in the magnetic field in the cluster and turbulence by the authors. Various geometries of merger have also been discussed. It is important to realise that the synchrotron spectrum is curved by nature (energy losses proportional to E$^2$). To understand the curvature in the spectra of different regions and identification of the break frequencies, spectral index maps between multiple frequencies are necessary. So far a study with two spectral index maps has been carried out for the cluster A2255 by Pizzo \\& Bruyn (2009). Radio halos and relics are rare sources and thus only a few clusters where such detailed study can be carried out are available. \\par One such cluster is Abell 2256 (hereafter A2256) which is a host to a radio halo and a radio relic and shows clear signatures of merger. Using two spectral index maps (150-350 MHz and 350-1369 MHz), a study of the complex dynamics in the cluster A2556 is presented here. A2256 is a rich, X-ray luminous ($L_{X [0.1 - 2.4 keV]}\\sim3.8\\times10^{44}$ erg s$^{-1}$, Ebeling et al 1996) galaxy cluster at a redshift of 0.0581 (Struble \\& Rood 1999). The X-ray surface brightness is elongated in the east-west direction and shows substructures (Briel et al 1991; Sun et al 2002); the radial velocity distribution of galaxies shows the presence of three distinct groups of galaxies (Berrington et al 2002). These have been interpreted to be indicators of ongoing merger in A2256. Recent temperature maps of A2256 with $Chandra$ (Sun et al 2002) and XMM Newton (Bourdin \\& Mazzotta 2008) show variation between 4 - 10 keV in $\\sim0.6$ Mpc region around the cluster centre. Apart from several head-tail radio galaxies, A2256 hosts diffuse radio emission in the north-west of the cluster centre (the radio relic) and at the centre (the radio halo). It has been studied in radio wavelengths for the past three decades (Bridle \\& Fomalont 1976; Bridle et al 1979; Rottgering et al 1994 (hereafter R94); Clarke \\& Ensslin 2006 (hereafter CE06); Brentjens 2008 (hereafter B08)). Polarization of $\\sim20-40\\%$ was detected at 1400 MHz in the relic by CE06; at 350 MHz it is unpolarized ($<1\\%$) (B08). A2256 also hosts the peculiar steep spectrum source 'F'; the optical identification of it is still being debated (B08). New steep spectrum sources have been detected in the 330 MHz images of A2256 made using the GMRT by van Weeren et al (2009); these are at the periphery of the cluster and are unrelated to the halo and the relic that will be discussed in this paper. \\par Study of the integrated spectrum (including the radio galaxies, the compact radio sources as well as the diffuse radio emission) of A2256 has shown that it steepens at low frequencies (B08); a property unique to this cluster. The properties of the relic in A2256 are inconsistent with the scenario of acceleration in structure formation accretion shocks (Ensslin et al 1998). They propose a shock radius of $\\sim1$ Mpc for the geometry of the relic in A2256, which implies an origin in shocks interior to the cluster -- the merger shocks. Based on the simulations of Roettiger et al (1995), merger between clusters having mass ratios 2:1, in the direction northwest to southeast, such that the smaller cluster is moving towards the observer has been discussed in the case of A2256. This model was proposed to explain the X-ray properties and temperature distribution in A2256 as estimated by {\\it ROSAT} (Briel \\& Henry 1994); further sensitive X-ray measurements ({\\it ASCA, Chandra, XMM Newton}) do not confirm the measurements of temperature by the {\\it ROSAT} (Sun et al 2002; Bourdin \\& Mazzotta 2008); rendering the model unusable. Based on X-ray substructure and identification of three distinct groups of galaxies in A2256, two mergers have been discussed (Sun et al 2002; Berrington et al 2002; Miller et al 2003). Berrington et al (2002) propose two mergers -- one between two comparable mass subclusters and another between a group and the primary cluster. Miller et al (2003) have favoured the possibility of the group being responsible for the radio relic whereas CE06 favour the merger of two subclusters. The case of A2256 is complex.  The spectral distribution across the diffuse radio emission in A2256 is presented in this paper and its implications to the mergers are discussed. \\par This paper is organized as follows. In Section 2, radio observations of A2256 using the Giant Metrewave Radio Telescope (GMRT) and the data reduction are described. The radio image at 150 MHz and the spectral index maps between 150-350 MHz and 350-1369 MHz are presented in Section 3. Discussion of these results in the contexts of temperature of the ICM and the dynamics in the cluster are presented in Section 4. The conclusions are presented in Section 5. \\par We adopt a $\\Lambda$-dominated cosmology with a Hubble constant, $H_0$, of 73 km s$^{-1}$ Mpc$^{-1}$,  $\\Omega_M=0.27$ and $\\Omega_{\\Lambda}=0.73 $  which implies a distance scale of $\\sim 70$ kpc arcmin$^{-1}$ at the redshift of A2256. ",
        "conclusions": "We have carried out a multi-wavelength analysis of the merging rich cluster of galaxies A2256. Using new radio observations at 150 MHz from the GMRT and archival observations from the VLA (1369 MHz) and the  WSRT (350 MHz), we have produced spectral index images of the diffuse radio emission in A2256 over the range 150-1369 MHz. These spectral index images show regions of the diffuse radio emission having flat spectral indices in the NW and having steep spectral indices in the SE of the cluster centre. The implied synchrotron life times for the relativistic plasmas are in the range 0.08 - 0.4 Gyr. Such a distribution of spectral indices is interpreted as resulting from a merger through the cluster from the SE to the NW in the last 0.5 Gyr or so. Acceleration due to shocks can explain the emission only in the NW relic region. The generation of the diffuse radio halo emission is likely to be due to the turbulence injected in the ICM by the mergers. The injection of fluid turbulence by mergers in the last Gyr and the timescales of $\\sim$Gyr required for the cascade of MHD turbulence are consistent with the spectral ages $\\sim0.4$ Gyr. Furthermore, the diffuse radio emission shows spectral steepening toward lower frequencies. This low frequency spectral steepening is consistent with a combination of spectra from two populations of relativistic electrons created at two epochs (two mergers) within the last $\\sim$0.5 Gyr. Earlier interpretations of X-ray and optical data suggested that there have been two mergers in Abell 2256 in the last $\\sim$0.5 Gyr, consistent with the current findings. Also highlighted is the futility of correlating the average temperatures of thermal gas with the average spectral indices of diffuse radio emission in respective clusters. Spectral index imaging of diffuse radio emission in galaxy clusters is found to be a powerful tool to study cluster mergers, dynamics, and history of cluster formation. \\begin{figure} \\epsscale{0.8} \\plotone{fig1.ps} \\caption{A2256: GMRT 150 MHz image with a synthesised beam of $67''\\times67''$ and an rms of $3.6$ mJy beam$^{-1}$ . Contours are at -18, 18, 36, 54, 72, 90, 126, 180, 240, 360, 480, 600 mJy beam$^{-1}$. The DSS R band image is shown in greyscale. The linear scale is marked for an assumed redshift of 0.058. The labels A, B, C, D, F, O, G and H for the radio sources are according to the convention in Bridle et al (1979) and R94. \\label{fig1}} \\end{figure} \\begin{figure}[h] \\epsscale{1} \\plotone{fig2.ps} \\caption{Grayscale (colour in the online version) represents the spectral index of radio emission in A2256 between 350 and 1369 MHz({\\it left}) and between 150 and 350 MHz ({\\it right}). The synthesized beam is $67''\\times67''$. The symbols square, circle and star (left panel) represent the centroids of the primary cluster, the sub-cluster and the group, respectively (Berrington et al 2002). In the left panel, the black (white in the online version) contours are a schematic representation of regions at different temperatures (Bourdin \\& Mazzotta 2008). Contours enclose regions with temperatures of 4-6 keV (innermost), 6-8 keV (region surrounding the innermost region) and 8-10 keV (region toward SE) respectively. \\label{fig2}} \\end{figure} \\begin{figure}[h] \\epsscale{1} \\plottwo{fig3a.ps}{fig3b.ps} \\caption{Spectral index error maps in greyscale. {\\it Left} 350-1369 MHz. {\\it Right} 150-350 MHz. Overlaid are the images at 350 and 1369 MHz in the left and the right panels respectively. Synthesized beams in both the panels are $67''\\times67''$. Contours (starting at $5\\sigma$) are at: {\\it (Left)} 3.0, 4.8, 6.6, 9.0, 18.0, 36.0, 72.0, 144.0 mJybeam$^{-1}$ and {\\it (Right)} 0.35, 0.56, 0.77, 1.05, 2.1, 4.2, 8.4, 16.8 mJybeam$^{-1}$. \\label{fig3}} \\end{figure} \\begin{figure}[h] \\epsscale{0.4} \\plotone{fig4.ps} \\caption{Spectra of diffuse radio emission from six patches spaced equally along the broken line shown in Fig. 2 (right panel). The spectral indices, ($\\alpha_{150}^{350}$, $\\alpha_{350}^{1369}$), for the spectra from top to bottom are: (-0.60, -0.72), (-1.03, -0.59), (-1.00, -0.36), (-1.25, -0.59), (-1.85, -0.97) and (-2.30, -1.23). Typical errors are indicated in the second spectrum from the top for which there is no scaling in the flux density. All other spectra are shifted along the y-axis. \\label{fig4}} \\end{figure}"
    },
    "1005/1005.3574_arXiv.txt": {
        "abstract": "The (778)~Theobalda asteroid family attracted little attention so far, but our study shows that it is important in several aspects. In this paper we investigate the origin and evolution of Theobalda family. Firstly, we identify the family as a statistically relevant group in the space of synthetic proper elements. Using the hierarchical clustering method and adopted cut-off velocity of $d_{cutoff}=85$~ms$^{-1}$ we found that Theobalda family currently consists of 128 members. This family is located in the outer belt, near proper semi-major axis $a_{p}$~$\\approx$~3.175~au. This region is crossed by several three-body mean motion resonances which give rise to significant chaotic zones. Consequently, the majority of family members reside on chaotic orbits. Using two independent methods, chaotic chronology and backward integration, we found Theobalda family to be only 6.9~$\\pm$~2.3~Myr old. We have also estimated, that the family was likely produced by the cratering impact on a parent body of diameter $D_{PB}$~$\\approx78\\pm9$~km. ",
        "introduction": "\\label{s:intro}  Asteroid families are believed to originate by a catastrophic disruption of large asteroids. To identify an asteroid family, one looks for clusters of asteroids in the space of proper orbital elements \\citep{MilKne90,MilKne94}: the proper semi-major axis ($a_{p}$), proper eccentricity ($e_{p}$) and proper inclination ($I_{p}$). The orbital elements describe the size, shape and tilt of orbits. Proper orbital elements, being more constant over time than instantaneous ones, provide a dynamical criterion of whether or not a group of bodies has a common ancestor. Up to now, ejecta from a few tens of major collisions have been discovered in the main belt \\citep[e.g.][]{zappala94,Diniz05}.  The size and velocity distributions of the family members provide important constraints for testing our understanding of the break-up process, but erosion and dynamical evolution of the orbits over time can alter the original signature of the collision. It is nowadays well known that the kinematical structures of the asteroid families evolved over the time, with respect to the original post-impact situations, due to chaotic diffusion, gravitational and non-gravitational perturbations \\citep{MilFar94,bottke2001,carruba03,delloro04}. These mechanisms changed the original shapes of the families produced in collisions, and consequently complicated physical studies of high-velocity collisions.  Unfortunately, most of the observed asteroid families are old enough (older than 100 Myr \\citep{nes2006}) to be substantially eroded and dispersed. On the other hand, young asteroid families (younger than 10~Myr) such as Karin, Veritas and Iannini \\citep{nes2002,nes2003} or even very young families (younger than 1~Myr) such as Datura, Emilkowalski, 1992YC2, and Lucascavin \\citep{nesvok2006}, suffer little erosion during the period of time after a breakup event. Thus, they provide a unique opportunity to study a collisional outcome almost unaffected by orbit evolution.  In this paper we study Theobalda asteroid family. We present its basic properties including the identification of its membership, and the study of cumulative absolute magnitude distribution of the family members. Moreover, the diameter of the parent body has been estimated. We studied in detail the dynamical characteristics in the region occupied by the Theobalda asteroid family and analyzed the role of the dynamics in shaping the family .  As was noted by \\citet{nov09} this family is a very good candidate to estimate its age by the method of chaotic chronology (MCC). In order to apply MCC the family has to be located in the region of the main asteroid belt where diffusion takes place. Also, it is necessary that diffusion is fast enough to cause measurable effects, but slow enough so that most of the family members are still forming a robust family structure. As we show, it turns out that Theobalda family is an excellent case in this respect. Given that, as our main result we have estimated the age of the family. Using two different methods, MCC \\citep{menios07,nov09} and backward integrations \\citep{nes2002}, we estimate the age of the family to be 6.9~$\\pm$~2.3~Myr. Thus, we establish it as another young asteroid family.  The paper is organized as follows: In Section~\\ref{s:basic} we present the basic properties of the Theobalda family. We use the hierarchical clustering method (HCM), proposed by \\citet{zappala90}, to identify family members. Next, the cumulative absolute magnitude distribution of identified family members is discussed, and the size of parent body is estimated. The dynamical characteristics of the region occupied by Theobalda family are presented and discussed, and main mean motion and secular resonances, in that region, identified. The dynamical stability of the family members is analyzed, in particular the stability of the largest member of the family (778)~Theobalda. In Section~\\ref{s:age} we estimate the age of the family. This is performed firstly by using the backward integration method, and then by using the method of chaotic chronology. The good agreement between these two results indicates a reliable age determination. Finally, in Section~\\ref{s:conclusions} we summarize our results, discus some possibly interesting relations to other works, and draw our conclusions. ",
        "conclusions": "\\label{s:conclusions}  We have presented here a detailed study of Theobalda asteroid family. We found that family now consists of 128 members. By analyzing SFD of the identified family members we were able to infer diameter of the parent body to be $D_{PB}$~$\\approx78\\pm9$~km. However, this estimate is based on the assumption that all family members have the same albedo as the largest family member, asteroid (778)~Theobalda. In order to obtain better estimate, the albedos of as many as possible family members are desirable. Ongoing projects, such as Wide-Field Infrared Survey Explorer (WISE) should improve situation significantly in this respect.  The most, but not all, of Theobalda family members move on chaotic orbits, thus, giving rise to the significant chaotic diffusion which has been changing the kinematical structure of the family over time. The study of dynamical characteristics, in the region occupied by the family, showed that three three-body MMRs are the most efficient in shaping the family. These are (3,~3,~-2), (5,~-2,~-2) and (7,~-7,~-2) resonances, and they are connected in this high eccentricity region, allowing bodies to switch from one resonance to another.  The fact that some of the family members have stable orbits was the reason why we were able to apply backward integration method to estimate the age of the family. On the other hand, presence of the chaos in the region occupied by the family, allows to use MCC in order to estimate the age. Using both methods, and combining the results, we found the age of Theobalda family to be 6.9~$\\pm$~2.3~Myr. Given the very good agreement between results obtained with different methods as well as when applied to different groups, we believe this estimate is very robust. Thus, this is another family younger than 10~Myr. This result has several important implications, and some of them we mention bellow.  The young asteroid families are also known to be source of solar system dust bands \\citep[see e.g.][]{grun85,nes2003}. The origin of three main dust bands is known, and they correspond to Karin, Veritas and Beagle family \\citep{nes2006b,nes2008}. Also, the very young Emilkowalski family is the most probably source of incomplete dust band at $17^{o}$ \\citep{espy2009}. On the other hand, the origin of some less prominent bands, such as so called M/N dust band, is still not quite clear. The Theobalda family's young age, and its proper inclination of $I_{p}$~$\\approx14^{o}.3$ suggest that it might be a possible source of M/N dust band ($I_{p}$~$\\approx15^{o}$) \\citep{sykes90}. On the other hand, this dust band was linked to (170) Maria asteroid family \\citep{reach97}, and more recently to (1521) Seinajoki cluster \\citep{nes2003}. In any case, dust band produced by such a young family, as Theobalda, should be observable. The size of its parent body also suggests that it should produce a prominent dust band. If this is not M/N dust band, then there must be another dust band which can be linked to this family. Alternatively, it should be explained why and how this dust band has disappeared.  Theobalda asteroid family is located very close to the region where three (out of four) so-called \\textit{main belt comets} (MBCs)\\footnote{The MBCs are bodies with asteroid-like dynamical properties but comet like physical properties \\citep{hsieh2004}. These are dynamically ordinary main-belt asteroids on which, probably, subsurface ice has recently been exposed e.g. because of a collision.} have been discovered \\citep[see e.g.][]{jewitt2009}. As was suggested by \\citet{hsieh2009}, it is possible that this kind of bodies can be found among the members of other young families, probably many of which waiting to be discovered. Being young and dominated by C-type asteroids, we believe Theobalda family is very good place to start."
    },
    "1005/1005.2051.txt": {
        "abstract": " ",
        "introduction": "Fundamental scalar particles have not been found in Nature yet. Many examples  have been postulated over the last forty years; the most illustrious  being the Higgs boson, a cornerstone of the Standard Model (SM), which may well be the first one to be discovered   at the Large Hadron Collider. The existence of scalar fields has  also been invoked  in order to explain the inflationary phase in the early universe\\cite{lyth}. In the same vein, the late time acceleration of the  expansion of the universe could result from the presence of dark energy scalar fields\\cite{Copeland:2006wr,caldwell,trodden,Brax:2009ae}. In all these cases, the coupling of  scalar particles to normal matter leads to new interactions and possibly variations of parameters such as masses and coupling constants \\cite{uzan}.  The physics of scalar fields coupled to matter depends crucially on the scalar field mass. For instance, the Higgs boson is expected to have a large mass ($m_h > 115$ GeV) compared to the Hubble scale today ($H_0 = 10^{-42}$ GeV) implying  that the new interaction mediated by the Higgs boson is of very short range. Moreover  the Higgs field has been frozen at the minimum of its potential since the early stages of the universe. On the other hand, several extensions of the Standard Model  predict the existence of light scalar fields. Some of these field might be moduli of string compactifications or the dilaton. Another class comprises  the  fields which can appear as a consequence of the spontaneous symmetry breaking of a global quasi-exact symmetry. These pseudo-Goldstone bosons may play a relevant role in explaining the recent stage of acceleration observed in cosmological data. In this case, the scalar fields are nearly massless as their mass should be of the order of the Hubble rate now. All these light scalar fields can couple to matter  and hence introduce new long range forces which would appear as a modification of gravity. The new interaction would also induce a non-conservation of the matter field stress-energy tensor, which can be interpreted as arising from  mass-varying matter fields. Limits on this type of interactions arise from SNIa and CMB data \\cite{DEInt}. As we will see below, alternatively one can choose to work in the so-called Jordan frame, where the matter field stress-energy tensor is conserved. In this case, Newton's constant varies and its variation since Big Bang Nucleosynthesis should not exceed some forty percent \\cite{uzan}.   Modifications of gravity should not be taken lightly. There are stringent limits on the couplings of scalar fields to normal baryonic matter arising from the search of a fifth force or deviations from the usual gravitational force at different scales, from the laboratory to the solar system.  In particular, the Cassini bound~\\cite{Bertotti:2003rm} implies that the coupling to matter should be less than $10^{-5}$. Several mechanisms have been proposed to alleviate this problem. The first one which applies to the dilaton in the strong string coupling limit is due to Damour and Polyakov \\cite{DP} and it asserts that the field is attracted by the zero coupling limit cosmologically. Another mechanism occurs in the Dvali-Gabadadze-Porrati model \\cite{DGP} whereby the brane bending mode, whose presence would lead to strong deviations from Newtonian gravity, is suppressed at short distance thanks to the Vainshtein mechanism. Another possibility  is to generate a mass term for the scalar field which  depends on the matter density of the environment. This is implemented in the so-called chameleon mechanism \\cite{Khoury:2003rn,braxkhoury}, where a coupling between a scalar field (the chameleon field) and matter arises from a conformal rescaling of the metric in the matter sector. The nonlinear interaction of the chameleon field generates a shielding, known as the thin-shell effect, that diminishes the modifications of gravity for dense and large enough bodies.  In this paper we are interested in modifications of General Relativity in chameleon models, focusing on the possible consequences for the growth and collapse of large scale structures in the universe. The growth of perturbations in chameleon models in the linear regime was studied in \\cite{BraxLinear}. Recently, there has been a new bout of  interest in studying the influence of modified gravity models on the formation of large scale structures. Indeed most of the aforementioned models of modified gravity show very little deviation from a pure $\\Lambda$CDM model since Big Bang Nucleosynthesis. Differences only emerge at the perturbative level in the way structures grow. In that respect it is of high interest to go beyond the linear level and study the nonlinear regime of structure formation. A semi-analytical tool that has been used to analyse the nonlinear stages of gravitational collapse is the so-called spherical collapse model \\cite{SphColl}.  The spherical collapse model has been recently used in models with a simple Yukawa-type modification of gravity \\cite{MartinoStabenauSheth}, in the so-called $f(R)$ models \\cite{SchmidtLimaOyaizuHuIII}, in braneworld cosmologies \\cite{SchmidtLimaHu} and in models which allow for dark energy fluctuations \\cite{DEFluc}. We will focus on the collapse of objects with a thin shell that arises in the context of chameleon models.  This paper is organized as follows. In section 2 we give a brief introduction to chameleon models and  discuss the modifications of gravity which occur in these models, both for geodesics and the growth of perturbations. In section 3 we solve the dynamical equation for the scalar field in a spherical configuration of matter overdensity, finding its profile and studying the effects of the existence of a thin shell which effectively shields the modifications of gravity. These profiles serve as templates for the chameleon behaviour during the spherical collapse. Section 4 presents an analytical discussion of the consequences of the chameleon models for the spherical collapse models of large scale structure formation. Phenomenological consequences are expounded in section 5 with a simple example before presenting our conclusions. ",
        "conclusions": "We have investigated the effects of a chameleon scalar field on the nonlinear stages of the gravitational collapse of a spherical overdensity region immersed in a FRW background. The gravitational potential is effectively changed by the presence of the new long range interaction. The modification depends on the profile of the scalar field. We studied the profile for the situation at hand, finding two very distinct classes of profiles depending on the size of the collapsing region. In particular, for large enough collapsing region, the nonlinear dynamics of the scalar field produces what is  known as a thin shell effect, where the modification of gravity is essentially shielded. This has an influence on the dynamics of the shell as well as a potential effect on the critical density contrast for collapse, which in this regime of chameleon theories becomes scale dependent.  We have also studied the spherical collapse in the case of the original chameleon model with an inverse power law potential and a constant coupling to Cold Dark Matter. The effects are larger for the smaller collapsing regions. In all cases, the initial thin shell disappears quickly and the collapse proceeds with no shell until a potential thin shell reappears after virialisation. More general models such as $f(R)$ theories may lead to a stronger influence of the thin shell regime on the spherical collapse.   We have emphasized that an inhomogeneous matter density in the thin shell will certainly result from the influence of the chameleon of the spherical dynamics. The phenomenological results in this paper have been obtained neglecting this effect. Although we expect that this phenomenon will be negligible for large enough initial radii, for smaller radii we  expect that singularities such as caustics may appear. A more detailed analysis of these effects is left for future work."
    },
    "1005/1005.1571_arXiv.txt": {
        "abstract": "A new catalog of isolated galaxies from The Sloan Digital Sky Survey (DR5) is presented. 1520 isolated galaxies were found in 1.4 steradians of sky. The selection criteria in this so called UNAM-KIAS catalog was implemented from a variation on the criteria developed by Karachentseva 1973 including full redshift information. Through an image processing pipeline that takes advantage from the high resolution ($\\sim$ 0.4 $\\arcsec$/pix) and high dynamic range of the SDSS images,  a uniform $g$ band morphological classification  for all these galaxies is presented.  We identify 80\\% (SaSm) spirals (50\\% later than Sbc types) on one hand, and a scarce population of early-type E(6.5\\%) and S0(8\\%) galaxies amounting to 14.5\\% on the other hand. This magnitude-limited catalog is $\\sim 80\\%$ complete at 16.5, 15.6, 15.0, 14.6 and 14.4 magnitudes in the $ugriz$ bands respectively. Some representative physical properties including SDSS magnitudes and color distributions, color-color diagrams, absolute magnitude-color, and concentration-color diagrams as a function of morphological type are presented.   The UNAM-KIAS Morphological Atlas is also released along with this paper. For each galaxy of type later than Sa, a mosaic is presented that includes:  (1) a $g$-band logarithmic image, (2) a $g$ band filtered-enhanced image where a Gaussian kernel of various sizes was applied and (3) an RGB color image from the SDSS database. For E/S0/Sa galaxies, in addition to the images in (1), (2) and (3), plots of $r$ band surface brightness and  geometric profiles (ellipticity $\\epsilon$, Position Angle $PA$ and $A_{4}/B_{4}$ coefficients of the Fourier series expansions of deviations of a pure ellipse) are provided  \\footnote{The mosaic images are available via our on-line web-site \\\\ http://www.astroscu.unam.mx/$\\sim$hector/KIAS/kias.html}.    The size of the sample, the redshift completeness, the availability of high quality multicolor photometric data and detailed morphological and spectroscopic information  make the UNAM-KIAS catalog of isolated galaxies a suitable sample to address important issues such as: (1) comparative studies of environmental effects, (2)  constraining the currently competing scenarios of galaxy formation and evolution, (3) the nature and  evolution of elliptical and spiral galaxies in the field, (4) the spectral properties of a statistically significant number of isolated galaxies and their evolution as a function of redshift and (5) the fraction of AGN in isolated environments among other important topics.  The optimization and estimation of new structural parameters as well as important information to complement the existing one in other wavelengths is being carried out. ",
        "introduction": "\\label{S1}   The large and homogeneous redshift and image surveys like the 2dFGRS (Colles et al. 2001), SDSS (York et al. 2000) and more recently the 6dFGRS (Jones et al. 2004; 2009) have helped us to extend our view of the local universe. Not only the small-scale distribution of galaxies is being revealed but equally important, now it is possible to make accurate measurements of the relationship between the various physical properties of galaxies and the local environment. Recently, Park et al. (2007, 2008) and Park \\& Choi (2008) used the SDSS data and studied the environmental dependence of the observed morphology at large and small smoothing scales to address the question of whether galaxy morphology depends primarily on the large-scale environment in which the galaxy initially formed or on a smaller scale environment that may reflect the influence of later evolutionary effects such as galaxy-galaxy interactions.  Park et al. (2008) pointed out that galaxies statistically become more isolated if they recently merged and that at a fixed large-scale density, more isolated galaxies are more likely to be recent merger products.   Galaxy formation models within the $\\Lambda CDM$ paradigm are every time better tuned to fit the joint luminosity/color/morphology distribution of low-redshift galaxies (c.f. De Lucia et al. 2006). These models assume that bulge formation takes place during mergers and that a galaxy can grow a disk later on provided that it is fed by an appreciable cooling flow. Spheroid formation is extended by assuming that bulges can also grow from disc instabilities. This additional mechanism can change the relative fractions of morphological types in these models. In the simulations of  De Lucia et al.  (2006), the final fractions of ellipticals, spirals and lenticulars brighter than -18 in the $V$ band are  17\\%, 65\\% and 18\\%, very close to the fractions of 13\\%, 67\\% and 20\\% reported by Loveday (1996). If bulge growth is switched-off, the above fractions become 7\\%, 84\\% and 8\\% respectively, very close to the reported fractions for isolated galaxies from the CIG catalog (Karachentseva 1973) that are in common to the SDSS (DR6) by Hern\\'andez-Toledo et al. (2008). This could be evidencing the impact of secular processes in the properties of isolated galaxies.   Uniformly selected and observed samples of galaxies that have not suffered any interaction with another normal galaxy or with a group environment over a Hubble time are thus crucial for studying intrinsic and secular processes able to affect  the structure, morphology and dynamics of galaxies, for instance, the formation and evolution of bars, rings, lopsidedness, and bulges.  Homogeneous observational data for such isolated galaxies are crucial for obtaining transparent scaling relationships and correlations that can be appropriately confronted to model predictions (see, e.g. Zavala et al. 2003; Avila-Reese et al. 2008).  It is worth to mention that an important requirement for all these samples is that of a well defined and strong isolation criteria, along with a uniform quality data acquirement in several wavelengths, and of statistical completeness.   In order to overcome some limitations of one of the most representative census of isolated galaxies in the local Universe started a few decades back by Karachentseva (1973) (e.g. lack of uniform, deep and high resolution CCD imaging, scarce of uniform and reliable photometric data as well as incompleteness in high resolution spectroscopic data),  we compiled a new catalog of isolated galaxies  from The Sloan Digital Sky Survey (DR5) as part of the IAUNAM-KIAS collaboration. This hereafter called UNAM-KIAS catalog has been compiled  by implementing a variation on the criteria developed by Karachentseva, including full redshift information as well as more accurate photometric information now available in the SDSS database. Using the three dimensional information, we are able to avoid missing galaxies that happen to lie close to background or foreground galaxies. One principal aim of this new catalog is, providing new isolated galaxy candidates with the highest possible statistical and redshift completeness to complement the existing lists of similarly selected galaxies in the local Universe. Another relevant aim is to provide a body of observational parameters that will be used for a comparison analysis of the predicted properties of different models of galaxy formation and evolution.   In a companion paper by Hern\\'andez-Toledo et al. 2008 (hereafter Paper 1) the results of a detailed morphological reevaluation of all  the available isolated CIG galaxies (549) from the Karachentseva (1973) catalog in common to the SDSS (DR6) was presented. The methodology was based on an image processing scheme that takes advantage of the uniformity, photometric conditions and high resolution  (0.4 $\\arcsec$/pix) of the SDSS CCD images.  Here we have applied the same image procedures to proceed with a morphological classification for the isolated galaxies of the UNAM-KIAS catalog.  A good understanding of these isolated, non--perturbed, galaxies and in particular of their morphological content is at the basis of any further correlations with other physical parameters and is one of the main legacies in the presentation of the UNAM-KIAS catalog.   The outline of the paper is as follows. Section 2 summarizes the relevant SDSS information and selection criteria used to compile the UNAM-KIAS catalog of isolated galaxies. In Section 3 the results of a detailed morphological classification are presented along with the public release of the UNAM-KIAS morphological atlas. In Section 4 some general properties, namely, completeness, apparent magnitude, color distributions, color-color diagrams, absolute magnitude-color and concentration-color diagrams as a function of morphological type are presented. Finally, in Section 5 a brief summary of the principal results and the main conclusions achieved are presented. The corresponding photometric and spectroscopic catalogs will be presented elsewhere.  The widely accepted cosmology with $H_o=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m}=0.27$ and $\\Omega_{\\Lambda}=0.73$, suggested by the WMAP data (Spergel et al. 2003), is adopted throughout this paper. ",
        "conclusions": "  a) The spatial distribution of the UNAM-KIAS galaxies, as seen from their redshift distribution does not show evidence of obvious concentrations associated with major components of large-scale structure in the local Universe but a more careful study in this direction needs to be done.  This apparent homogeneity points to a sample that is close to a local homogeneous component of the isolated galaxy distribution.  b) Although the UNAM-KIAS galaxies comprise a nearby sample, its redshift interval has increased with respect to that in the CIG catalog by about a factor of two, making it deeper and representative of the galaxies in this environment in the local Universe.   c) The UNAM-KIAS sample is also diverse in terms of  physical properties explored here. A scarcity of  nearby dwarf isolated E/S0 galaxies and a  tendency of  bright E/S0 galaxies being farther away than spiral galaxies is observed.  d) The UNAM-KIAS catalog of isolated galaxies is a magnitude-limited sample that is reasonably  complete ($\\sim$ 80$\\%$) up to 15.2 $r$ band apparent magnitudes, providing statistical significance to studies of these galaxies in isolated environments.  The finding of reasonably nearby isolated galaxies uniformly-selected and with detailed morphological information is of high relevance. The UNAM-KIAS catalog along with the most isolated galaxies of the CIG catalog that are in common to the SDSS will provide a unique database that can be used for several studies, including (i)  studies of the environmental effects on galaxies belonging to groups and clusters, and (ii)  for confronting with theoretical and model predictions of galaxy evolution.   Other relevant aims, derived from the present study, are to provide as much as possible a new body of observational parameters for these isolated galaxies in different wavelengths that will greatly enable to a major understanding of the nature of the galaxies in this environment."
    },
    "1005/1005.1657_arXiv.txt": {
        "abstract": "Assuming homogeneous isotropic $\\Lambda$CDM cosmology allows $\\Lambda$, spatial curvature and dark matter density to be inferred from large scale structure observations such as supernovae. The purpose of this paper is to extend this to allow observations to measure or constrain inhomogeneity and anisotropy. We obtain the general inhomogeneous anisotropic $\\Lambda$CDM solution which is locally asymptotic to an expanding de Sitter solution as a late time expansion using Starobinsky's method (analogous to the `holographic renormalization' technique in AdS/CFT) together with a resummation of the expansion. The dark matter is modeled as perfect dust fluid. The terms in the expansion systematically describe inhomogeneous and anisotropic deformations of an expanding FLRW solution, and are given as a spatial derivative expansion in terms of data characterizing the solution - a 3-metric and a perturbation of that 3-metric. Leading terms describe inhomogeneity and anisotropy on the scale set by the cosmological constant, approximately the horizon scale today. Higher terms in the expansion describe shorter scale variations. We compute the luminosity distance-redshift relation and argue that comparison with current and future observation would allow a partial reconstruction of the characterizing data. We also comment on smoothing these solutions noting that geometric flows (such as Ricci flow) applied to the characterizing data provide a canonical averaging method. ",
        "introduction": "Modern large scale structure observations, such as the SDSS low redshift studies of `standard candle' supernovae \\cite{Frieman:2007mr,Sako:2007ms}, afford us the capability to directly test our cosmological model assuming only the Einstein equations, matter content relevant for our current epoch, and homogeneity and isotropy (and hence the FLRW metric). Historically Hubble's measurement of the current expansion rate was the first test of this type. Recently the famous results of the Supernova Search Team \\cite{Riess:1998cb} and the Supernova Cosmology Project \\cite{Perlmutter:1998np} allowed comparison of luminosity distance-redshift for various matter content models with supernovae data up to redshift $\\sim 5$, showing the likely current domination of a positive cosmological constant over dark matter.  An important point is that unlike CMB observations, large scale structure allows inferences to be made about the parameters and dynamics of a model without any assumption on initial conditions, except for the, albeit strong, assumptions of homogeneity and isotropy. While the cosmological constant can be estimated from CMB measurements, this inference is far less direct, and is subject to a choice of inflationary initial conditions. Here we are certainly not arguing that one should not take inflationary initial conditions seriously, but rather we wish to emphasize that when one makes inferences about physical models from observational data, the strength of the parameter estimation implicitly depends on the strength of the assumptions made. It is useful to have a variety of methods to allow estimation with differing assumptions.  It is worth emphasizing that in the context of inflation it is the inflationary mechanism that suppresses large scale inhomogeneity, anisotropy as well as spatial curvature. Thus if one is interested in assessing the magnitude of spatial curvature by measurement from observation, for example if one thought inflation lasted only a minimal time and therefore interesting structure may be seen on large scales (e.g. \\cite{Efstathiou:2003hk}), it is conceptually inconsistent to include spatial curvature. One should consider general anisotropic inhomogeneous deformations at the same time as introducing spatial curvature, unless one supplements the initial conditions before inflation with homogeneity and isotropy, as for example occurs after semiclassical bubble nucleation of an open universe \\cite{Coleman:1980aw}.  The purpose of this paper is to allow the comparison of a $\\Lambda$CDM cosmology with large scale structure data such as supernovae, \\emph{without making the assumptions of homogeneity or isotropy} on large scales. This is not because we are arguing the universe is far from homogeneous or isotropic, but simply because in order to argue the universe is homogeneous and isotropic, ideally one should deduce this from the data without assuming it. Furthermore with future measurements one would hope to be able to quantify precisely \\emph{how far} from homogeneity and isotropy the universe is, or at least put rigorous upper bounds on how large these quantities are on different scales.  Of course in order to do any of this, one must be able to compute and characterize general inhomogeneous solutions for a $\\Lambda$CDM cosmology. In addition, one must then be able to compute quantities relevant for observation, such as the luminosity distance-redshift relation, and show how this allows inferences about the data that characterizes the solutions to be made from real large scale structure data. Certain exact solutions have allowed such analysis to be performed for very special cases, for instance using the LTB solution which is isotropic but inhomogeneous, or the Bianchi models which are anisotropic but homogeneous \\cite{Lemaitre:1933gd,Tolman:1934za,Bondi:1947av,Mustapha:1997xb,Mustapha:1998jb,Celerier:1999hp,Chung:2006xh}, but general exact solutions are likely not to exist in any closed form.  In this paper we show how to construct general inhomogeneous, anisotropic solutions of the Einstein equations with cosmological constant and dust fluid as a model for baryonic and dark matter using a method which we call a `late time expansion'. This method was first developed by Starobinsky \\cite{Starobinsky:1982mr} in the context of inflation, and is similar to methods employed in AdS/CFT holographic renormalization (see \\cite{Skenderis:2002wp}). Rather than starting from the conventional approach of cosmological perturbation theory about the FLRW solution, instead we use the fact that provided the local universe is dynamically dominated by the cosmological constant, one can describe the solution in a precise way as a deformation about de Sitter space, characterized by its late time asymptotic behaviour, or what we shall term the \\emph{final} boundary conditions. The leading late time behaviour of the solution is a universe that is locally de Sitter. Choosing the late time de Sitter conformal boundary to have proper time, $t \\rightarrow \\infty$, the full solution then takes the form of an expansion in positive powers of $e^{- \\Lambda t/3}$, where higher terms in the expansion give inhomogeneous corrections to this local de Sitter behaviour, and are given in terms of increasing numbers of derivatives acting on the final boundary data. We find the form of this final data is given by a 3-metric and a perturbation of that 3-metric. The leading terms in the late time expansion describe inhomogeneity and anisotropy on the scale set by the cosmological constant, and hence for our universe would correspond to approximately today's horizon size deformations. Higher terms in the series then govern derivative corrections which correspond to increasingly smaller scale inhomogeneity and anisotropy.  The first part of our paper is devoted to giving the solution for $\\Lambda$CDM, taking dust fluid, using the `late time expansion' about an asymptotically locally de Sitter solution with flat spatial slices. We give results beyond leading order, and in particular discuss the final data that arises in the expansion and characterizes the solution. For dust fluid it is possible to take a normal coordinate system adapted to the conformal boundary that is comoving with the dust - this is not possible for fluids with pressure. This simplifies the calculation considerably, and is what allows us to give results straightforwardly beyond the leading order.  In the second part of the paper, having showed how to give inhomogeneous anisotropic solutions as a late time expansion deforming about flat sliced de Sitter, we go on to consider the convergence of the late time expansion, arguing that whilst one might naively expect it to break down around $\\Lambda$-matter equality, it actually is far better, extending back to very early epochs. We then discuss the drawback of our first expansion, namely that homogeneous isotropic dust and curvature are also described as deformations and treated order by order in the late time expansion, even though we know the exact homogeneous isotropic solution is an expanding FRLW universe, whose solution we may treat analytically. However, we show that the result may simply be modified to give an expansion about the exact expanding FRLW solution using a simple but powerful series resummation on the late time expansion. The terms in the expansion then describe deformations only involving inhomogeneity and isotropy. If the deviations about FLRW solution are small we believe, although here do not prove, that this late time expansion converges back to high redshifts, certainly sufficient to describe the region relevant for large scale structure observations.  The third part of the paper is then concerned with applying these inhomogeneous and anisotropic solutions to observations of our universe. We compute the luminosity distance-redshift relation explicitly for the first two non-trivial orders in the expansion. We then discuss how one might in principle use this to compare to supernovae data and therefore extract information about the characterizing data, the boundary 3-metric and its perturbation. In particular we discuss how the boundary data is extracted as a derivative expansion about the point corresponding to the observers asymptotic position, and that fitting increasingly high orders in the luminosity distance-redshift relation expansion to data gives information about increasingly high derivatives of the local curvature of the 3-metric and perturbation. Thus information characterizing the inhomogeneous and anisotropic universe we live in can be systematically extracted from observational data, provided that, as we argue, we are living in and observing the regime where the late time expansion converges. In fact not all characterizing data can be extracted from this luminosity distance-redshift measurement, and we deduce the subset of the data that may be extracted.  We briefly conclude with a discussion of averaging or smoothing of our inhomogeneous solutions. Since in our method the late time inhomogeneous cosmologies are determined by geometric final data consisting of a 3-metric and its perturbation, geometric flows, such as Ricci flow, provide the canonical notion of averaging this characterizing data. Considering this we arrive at a natural averaging of our solutions similar in spirit although rather different in origin to that discussed by Buchert and Carfora \\cite{Buchert:2002ht,Buchert:2003}.  It is simple to argue that provided one begins with small deformations of FLRW, such smoothing flows to the FLRW solution. ",
        "conclusions": "Suppose one wishes to consider spatial curvature in a cosmological model as is often done. As we emphasized in the introduction, at least in the context of inflation and presumably more generally, it is unnatural to introduce spatial curvature without considering inhomogeneity and anisotropy. Once one introduces spatial curvature without assuming homogeneity and isotropy, one is typically in a regime beyond that described by cosmological perturbation theory. For example, as we have discussed above, if the spatial curvature is introduced as a squashed 3-sphere, rather than a round 3-sphere, then for any appreciable squashing, cosmological perturbation theory will fail as the solution is not a small metric deformation of FLRW. Thus if one wishes to observe or constrain the magnitude of spatial curvature in our universe, one is naturally led to the more general problem of characterizing curvature, inhomogeneity and anisotropy non-linearly in the amplitude of the metric and matter field perturbations. It is this characterization, and construction of the associated cosmological solution that has been the purpose of this paper.  Using Starobinsky's late time expansion (analogous to techniques developed for `holographic renormalization' in AdS/CFT) together with a resummation of the expansion, we are able to find the general inhomogeneous anisotropic solution to cosmological constant and dust fluid matter that is locally asymptotic to expanding de Sitter. The terms in the expansion describe a systematic deformation about a homogeneous isotropic expanding FLRW solution, and are arranged as a derivative expansion in the data that characterizes the solution.  The solution is fully non-linear in the amplitude of the metric and matter field deformations from FLRW. The leading terms in the expansion describe deformations on horizon scales, and higher terms describe deformations on increasingly small scales. The `final' data lives on the conformal boundary of the asymptotic de Sitter future. It is given by the boundary metric $\\bar{g}_{ij}$ and the tensor $h_{ij}$ that lives on that boundary geometry, and as we have discussed, this pair can be elegantly thought of as defining a 3-metric and a perturbation of that 3-metric (up to an inequality on the trace of the perturbation that must be satisfied to ensure positive dust density). Truncating the expansion at some finite order then allows a controlled approximation to the true  inhomogeneous anisotropic solution.  Recent cosmological observations indicate that we are resident in a universe whose matter content primarily consists of a cosmological constant and cold dark matter. We have argued that ignoring radiation is justified provided we restrict attention to the recent matter era ($z \\ll 10^3$). We expect that for weak inhomogeneity and isotropy, the series given by our late time expansion should converge for this recent matter era, and furthermore, the convergence will be fast. Thus truncating the series at some low finite order will provide a good and controlled approximation to the true solution. Hence we hope that the solution is more than simply a formal construction, and may actually allow us to interpret observations in light of our assumptions that the Einstein equations hold and that the late time matter content is $\\Lambda$ and dust fluid, without assuming anything about the homogeneity or isotropy of our universe. Ideally current and future structure observations, such as standard candle supernovae observations, could allow us to characterize systematically and precisely the large scale inhomogeneity and anisotropy, measuring them or constraining their size.  Our calculation of the luminosity distance-redshift relation to leading non-trivial orders in principle allows one to compare to large scale structure data and directly determine the Riemann tensor of the boundary metric and the value of the $h_{ij}$ tensor at the position corresponding to the observer position. These give the first anisotropic corrections to the solution. It allows one to extract partial information on the derivative of that curvature, giving $\\bar{\\nabla}_{(i} \\bar{R}_{jk)}$, but not the antisymmetric components. This represents the leading inhomogeneity in the solution. Working to higher order in the expansion, one would extract more information about the future boundary data $(\\bar{g}_{ij}, h_{ij})$. It would be interesting to understand whether there are parts of the data that are inaccessible by the luminosity distance-redshift measurement, and if so whether they might be determined by other types of observation, and precisely how much data can  be extracted in principle \\cite{Ellis}.  The use of an effective perfect fluid description of the dark matter already implies that an averaging has been performed. We consider averaging in the context of our inhomogeneous and anisotropic late time solutions. We argue that since the final data can be described in terms of a 3-geometry and its perturbation, it is natural to consider geometric flows, such as Ricci flows, to smooth the data. We consider Ricci flow of the boundary metric $\\bar{g}_{ij}$ which then induces an associated flow on perturbations of that metric, and hence on the tensor $h_{ij}$. We confirm that using Ricci flow smoothing, a small perturbation away from flat sliced FLRW is smoothed towards the FLRW solution as we would wish."
    },
    "1005/1005.2524_arXiv.txt": {
        "abstract": "Nuclear Stellar Disks (NSDs), of a few tens to hundreds of parsec across, are a common and yet poorly studied feature of early-type galaxies. Still, such small disks represent a powerful tool to constrain the assembling history of galaxies, since they can be used to trace to the epoch when galaxies experienced their last major merger event. By studying the fraction and stellar age of NSDs it is thus possible to test the predictions for the assembly history of early-type galaxies according the current hierarchical paradigm for galaxy formation. In this paper we have produced the most comprehensive census of NSDs in nearby early-type galaxies by searching for such disks in objects within 100 Mpc and by using archival images from the Hubble Space Telescope. We found that NSDs are present in approximately 20\\% of early-type galaxies, and that the fraction of galaxies with NSDs does not depend on their Hubble type nor on their galactic environment, whereas the incidence of NSDs appears to decline in the most massive systems. Furthermore, we have separated the light contribution of twelve such disks from that of their surrounding stellar bulge in order to extract their physical properties. This doubles the number of decomposed NSDs and although the derived values for their central surface brightness and scale-length are consistent with previous studies they also give a hint of possible different characteristics due to different formation scenario between nuclear disks and other kinds of large galactic disks. ",
        "introduction": "\\subsection{Selection and Acquisition}  For this study we have selected, using the LEDA database \\citep{paturelet03}, all early-type galaxies ($E$, $E$-$S0$, $S0$ and $S0a$) within 100 Mpc so as to produce a volume-limited sample. This list was then cross-correlated with the HST archive and all the galaxies with images from the WFPC2 and ACS cameras were requested, preferably the filters F555W and F606W. In some cases these filters were not available and others had to be chosen. The use of different passbands should not bias our estimates of the disk inclination and extent, however, as shown by \\citet{mor10} nuclear disks do not present strong radial colour gradients.\\\\ After downloading the available images they were checked for their quality and those saturated, too dusty or in which the galaxy was lying at the edge of the detector were discarded since in these cases it would have been impossible to find NSDs even if they were present. In the end we collected a sample of 466 early-type galaxies with HST images, out of a parent LEDA sample of 6801 galaxies.  \\subsection{Sample Properties}  Having a sample, we retrieved the global parameters of our galaxies. Of particular interest are M$_{\\mathrm{B}}$ and an estimate of the galactic environment, because they directly relate to the general questions we are trying to answer, and the galaxy inclination because it influences our ability to find the disks.The magnitude of our galaxies and an estimate for their inclination values are available from the LEDA database whereas to obtain a value for the galactic density we used the Nearby Galaxies Catalogue of \\citet{tully88}.  Although the distribution of M$_{\\mathrm{B}}$ and environment values for our HST sample and its parent LEDA sample are somewhat different this does not pose as a problem when we are dealing with the incidence of NSD. Indeed in this case it only matters that in each magnitude bin we have enough objects to obtain secure estimates for the NSD fraction. Also inclination differences do not matter for constraining the incidence of NSDs, as long as the distribution of inclination in each HST subsample allows a reliable correction of the fraction of NSD (that is if there are enough galaxies where the disk could have been detected, see \\S4.1.1). On the other hand we need to keep in mind that any conclusion on the structural properties of the NSDs, such as their typical size or mass, will be specific to our HST sample.  \\subsection{Inclination}  The more face-on a disk is, the harder it is to identify it. \\citet{rw90} have studied this problem and concluded that for inclinations with $cos(i) > 0.6$ it is impossible to detect disky signatures in the galaxy isophotes. The inclination, $i$, is defined as being 0$^{\\circ}$ for face-on disks and 90$^{\\circ}$ if they are edge-on.  Although we do not know for sure the inclination of the nuclear disks, we will assume that the vast majority of them will have the same inclinations of their host galaxies, which is justified if we consider our sample galaxies as oblate axisymmetric spheroids. The LEDA database provides inclination estimates for the host galaxies based on the method applied by \\citet{heid72} and \\citet{bot83}. The inclination was calculated assuming all galaxies of a given Hubble type had an axis ratio equal to that observed in the flattest galaxy of that type. By construction, this method provides a conservative lower estimate for our galaxies' inclination, since intrinsically rounder galaxies in a given Hubble type would in reality have a higher inclination. Using such lower estimates to correct the fraction of NSD will lead to an upper limit on their incidence, whereas the uncorrected fraction provides a lower limit as if we were assuming that all galaxies where edge-on. ",
        "conclusions": "Nuclear stellar disks have long been regarded as common features in early-type galaxies, but a quantitative assessment of such statements was still lacking prior to this work. By performing the most extensive volume-limited census of NSDs to date we have now shown that between 13 to 23\\% of nearby early-type galaxies host such systems. The incidence of nuclear stellar disks appears to decline in the most massive systems, consistent with the expectations that dry mergers dominated the most recent history of such objects \\citep{ks09,kb03} and also with the finding that NSDs are exceedingly rare in brightest cluster galaxies \\citep[2\\%,][]{laine03}. On the other hand, it is less clear what may cause the decrease in the fraction of NSDs that is found also at the faint end of our sample.  Similarly, the lack of a dependence on galactic environment seems to contrast the first-order predictions of semi-analytical models (Fig.~2). In fact, according to these galaxies in clusters should have assembled much earlier and thus there should have been more time to regrow a nuclear disk. On the other hand, galaxies in cluster may not have ample reservoirs of cold gas at their disposal, in particular if orbiting within the hot-gas halo of the cluster, so that the early assembly of clusters does not make the presence of nuclear disks in cluster galaxies much more likely than in their field counterparts.  Using the well established technique of \\citet{scoben95} to disentangle the light contribution of embedded disks from that of their surrounding bulges, we have extracted the structural parameters of 12 out of the 63 NSDs that we have identified. By doubling the number of NSDs we have shown even more convincingly that NSDs, like other kinds of large galactic disks, follow a  correlation between the central surface brightness $\\mu^0_c$ and scale length $h$, although it would appear that NSDs obey a somewhat different relation with smaller scale lengths or fainter central surface brightness values. Considering that disk galaxies obey a very precise relation between their total luminosity and their maximum rotation velocity (the Tully-Fisher relation, \\citet{tully77}) that may ultimately derive from the fact that they form in the dark matter halos that surround and dominate the mass budget of galaxies \\citep{steinmetz99}, it may not be completely unexpected that NSDs, by forming in the central and bulge-dominated regions of galaxies, may have different structural and dynamical properties than that of their larger relatives. In fact, it is also possible that the structural properties of NSDs vary when dealing with different kinds of early-type galaxies such as disky and boxy systems \\citep{kb96}. Although the main reason for failing to isolate the disk component in the presence of a bulge of varying boxiness with radius (Sec 3.3) has mostly to do with an unclear a priori for the intrinsic $A_4$ profile that the bulge should have in the absence of a disk, we cannot exclude that in boxy ellipticals NSDs may have a different profile than exponential. Along the lines first suggested by \\citet{bs99}, NSDs in boxy ellipticals could be akin to the rapidly rotating components that in these objects are generally identified as kinematically decoupled cores, although in general such structure tend to extend over much larger scales than NSDs \\citep{m06}, in particular for massive ellipticals with the oldest stellar populations.  Finally, although for the majority of the NSDs we could not clearly separate the disk and bulge contributions to the central light of their host galaxies, we have nonetheless estimated their sizes using as a reference radius the distance from the centre where the diskyness of the central isophotes appears to peak. Such scale length estimate, coupled with the $\\mu^0_c - h$ relation that NSDs appear also to obey and the N-body simulations, will be used in future investigations to assess which of our sample NSDs are indeed fragile against major merging events, thus providing good targets for a spectroscopic follow up aimed at constraining their stellar ages and thus, the time of their last major merger."
    },
    "1005/1005.3591_arXiv.txt": {
        "abstract": "A web-based, interactive system for the remote processing of imaging data sets (\\textit{i.e.}, EUV, X-ray and microwave) and the automated interactive detection of wave and oscillatory phenomena in the solar atmosphere is presented. The system targets localised, but spatially resolved, phenomena, such as kink, sausage, and longitudinal propagating and standing waves. The system implements the methods of Periodmapping for pre-analysis, and Pixelised Wavelet Filtering for detailed analysis of the imaging data cubes. The system is implemented on the dedicated data processing server \\url{http://pwf.iszf.irk.ru}, which is situated at the Institute of Solar-Terrestrial Physics, Irkutsk, Russia. The input data in the \\textsf{.sav}, \\textsf{.fits} or \\textsf{.txt} formats can be submitted via the local and/or global network (the Internet). The output data can be in the \\textsf{png}, \\textsf{jpeg} and binary formats, on the user's request. The output data are periodmaps; narrowband amplitude, power, phase and correlation maps of the wave's sources at significant harmonics and in the chosen spectral intervals, and \\textsf{mpeg}-movies of their evolution.  The system was tested by the analysis throughout the EUV and microwave emission from the active region NOAA 10756 on 4 May 2005 observed with TRACE and the Nobeyama Radioheliograph. The similarity of the spatial localisation of three-minute propagating waves, near the footpoint of locally open magnetic field lines determined by the potential field extrapolation, in both the transition region and the corona was established. In the transition region the growth of the three-minute amplitude was found to be accompanied by the decrease in the line of sight angle to the wave propagation direction. ",
        "introduction": "\\label{Introduction} The detection and the classification of various solar features is a complex task that relies on the expertise of observers, modelers, and theorists. The identification of highly energetic features, such as flares and CMEs, is often straightforward, and their signatures are well known and usually uniformly accepted in the solar physics community. Also, large-scale features, such as coronal holes, active regions, and loops can be identified with standard image processing algorithms (\\textit{e.g.} Aschwanden \\textit{et al.}, 2008; Scholl and Habbal, 2008; Krista and Gallagher, 2009). The main difficulty lies in the identification of \\lq\\lq weak\" features, which are close to the instrumental resolution or signal-to-noise threshold. However, these features, such as waves, often can provide us with invaluable information on the physical conditions in the solar atmosphere (\\textit{e.g.} see the review by Nakariakov and Verwichte, 2005).  The study of coronal waves and oscillations is one of the most rapidly developing areas of solar physics and lies in the mainstream of the scientific exploitation of solar mission data.  Coronal-wave processes attract our attention because of their possible role in the enigmatic problems of coronal physics: coronal heating, triggering of flares and CMEs, and the acceleration of the solar wind. Also, as properties of waves contain information about the medium of their propagation, the waves are natural probes of the medium, providing us with a tool for its diagnostics. The analysis of coronal data in order to determine parameters of waves and oscillations is a non-trivial task, mainly because of the detection of these phenomena near the very threshold of the instrumental detectability, and the necessity to analyse the spatial, temporal, and phase variations simultaneously at the highest-possible resolution. The problem has been amplified by the lack of an instrument dedicated to coronal-wave studies. The \\textit{Solar Dynamics Observatory/Atmospheric Imaging Assembly} (SDO/AIA) is the first instrument to have coronal seismology as one of its scientific themes. AIA will provide us with greatly enhanced cadence and signal-to-noise ratio in comparison with the previous generation of solar-coronal imagers. In the past, the detection of waves in the solar corona relied heavily on individual examination of time sequences of images (\\lq\\lq movies\") by experts in the field. The availability of AIA will substantially increase the data rates, requiring the development of novel data analysis methods for the automated detection of coronal wave and oscillatory processes. The pioneering studies of this problem (De Moortel and McAteer, 2004; Nakariakov and King, 2007; Marsh \\textit{et al.}, 2008; McIntosh \\textit{et al.}, 2008) have addressed this issue, investigating automated-detection techniques of coronal waves, based upon the wavelet and Fourier transforms, Bayesian statistics and coherence/travel time based approach.  There are several classes of MHD waves in the solar corona, which are to be recognized by an automated detection system. The observational evidence of the presence of MHD waves and oscillations is abundant (see \\textit{e.g.} Nakariakov and Verwichte, 2005). Observed properties of these waves are determined by the original MHD mode (slow or fast magnetoacoustic, Alfv\\'en) and the guiding plasma structure. Different MHD modes of plasma structures have very different physical properties (\\textit{e.g.} compressibility, polarization, dispersion, cut-offs, \\textit{etc.}) and hence have different observational manifestations.  SDO/AIA will be able to see at least two of them, the transverse (kink) and the longitudinal modes at higher spatial and temporal resolution than previously.  In this paper we describe a web-based system for the remote processing of imaging datasets (such as SOHO/EIT, TRACE, STEREO/EUVI, Nobeyama Radioheliograph, \\textit{Hinode}/XRT, CORONAS-PHOTON/TESIS, and SDO/AIA, designed for the automated detection of periodic and quasi-periodic processes. The system has a modular structure that allows for the easy update and inclusion of new data analysis and visualization tools. The present version of the system implements the techniques of periodmapping (Nakariakov and King, 2007) and the Pixelised Wavelet Filtering (PWF: Sych and Nakariakov, 2008). Also, there is a possibility of using the standard Morlet wavelet and periodogram analysis of 1D signals.  The method of peridomapping provides fast preliminary analysis of imaging datacubes, allowing for the identification of the candidate regions and time intervals for the presence of spatially resolved oscillatory phenomena. With this method, the time is divided into 30 minutes (or, optionally, of other duration) time slots which overlap by a half of the duration. Periodograms of the time signal of each pixel (or optionally a macropixel, consisting of several adjacent pixels for improved signal-to-noise ratio) are constructed. Options include signal smoothing, user-defined filtering, and construction of the running-difference signal. Then, the frequency of the highest spectral power is identified and tested against a certain criterion. The default  is that if the highest spectral power is higher than five times the mean level of the spectral power, then the signal is treated as an \\lq\\lq oscillatory\" and the frequency of the highest spectral power is determined. Then, the pixel is assigned  a certain colour according to a chosen colour coding scheme, in which different colours correspond to different spectral frequencies. If the highest spectral power is lower than the detection threshold, the signal is treated as \\lq\\lq non-oscillatory\" and the pixel is left blank. Hence, a static 2D map of the analysed field of view (a \\lq\\lq periodmap\") for the analysed time interval is created. The blobs of the same or similar colour (corresponding to the same or similar time frequency of the oscillation) indicate the presence of the oscillation. Then the process is repeated for the next time interval. As a result, a time sequence of periodmaps is obtained, which can be very easily analysed by visual inspection, or automatically, identifying the \\lq\\lq oscillatory regions\" in certain time intervals.  The PWF method (Sych and Nakariakov, 2008) is a generalisation of the wavelet transform of 3D datacubes. Temporal signal of each spatial pixel is wavelet transformed (\\textit{e.g.} with a Morlet mother function), which results in the power, amplitude, and phase 4D data cubes (two spatial dimensions, time, and the frequency). The 4D data cubes can be processed according to a specific request. For example, selecting a certain spectral component or integrating over a certain narrow spectral range makes a 3D narrowband datacube that consists of a sequence of narrowband maps. The temporal evolution of the sources of the narrowband oscillations can be visualised by a movie (\\textit{e.g.} in the MPEG format). Similarly, one can create a movie of the global wavelet spectrum evolution. Selecting a certain 1D slit on the map makes it possible to make narrowband or broadband time--distance plots. In particular, this allows one to identify the observed mode, to discriminate between standing and propagating modes, determine the morphology of wave sources and study its time evolution, and detect multiple periodicities and their interrelation. This knowledge is crucial for MHD coronal seismology and the study of the solar atmospheric connectivity. In the future, we plan to include other data processing tools, such as the Gradient Pattern Analysis (GPA-method: Rosa \\textit{et al.}, 2000).  The main aim of the present system is the remote processing of data cubes supplied by the user.  The user does not need to have deep knowledge of data processing techniques, and programming. Also, it is not necessary to have the Interactive Data Language (IDL) installed on the user's computer. The minimal prerequisites are access to the Internet and a standard web browser. The input parameters are the imaging data cubes. The output parameters are the global and local (\\textit{e.g.} for a selected pixel or macropixel, or a slit) spectra, narrowband images, periodmaps, wave vector maps, and movies showing the temporal evolution of the maps.  Application of the system to the detection of coronal waves and oscillations is illustrated with the use of the microwave and EUV imaging data cubes of NOAA 10756 obtained with TRACE and the Nobeyama Radioheliograph (NoRH). During the passage of this active region across the solar disk on 1\\,--\\,5 May 2005 both standing and propagating waves have been detected (see Sych \\textit{et al.}, 2009 for details), which made it a good test case for the present study.  The paper is organised as follows. In the next section we describe some existing examples of on-line data access and analysis tools. Section 3 describes the scientific objectives of the system. Section 4 discusses the system architecture and gives an example of the user interface. Sections 5 and 6 illustrate the application of the system to the analysis of observational data from the spatial and temporal points of view, respectively. The conclusions and discussion of the results obtained are given in Section 7. ",
        "conclusions": "The principal aim of this paper is presentation of the specialised computational resource \\url{http://pwf.iszf.irk.ru} for the automated detection and analysis of wave and oscillatory processes in the solar atmosphere by remote access, and the demonstration of its efficiency. The computational tool presented has a modular structure, which includes algorithms for the automated localisation of oscillatory processes in time and space, analysis of their properties, and visualization of the results. The system allows for easy change of the modules and parallel execution, which is crucial for fast performance and for possible future applications for real-time data analysis. The system is based upon the periodmapping technique for the preliminary analysis and identification of the region and times of interest, and the PWF technique for the detailed analysis. The system provides for the determination of the spatial-time-frequency-amplitude-phase dynamics and structure of the waves. In particular, the system allows distinguishing between standing (with respect to the LOS) and propagating waves. From the user's point of view, the system is a web site that can be interacted with through a standard Internet browser and does not require installation of any specific software. We are not aware of the existence of analogous resources in solar physics, or in other areas of science and technology.  The efficiency of the discussed system was tested by the analysis of the three-minute waves over the sunspot NOAA 10756 observed in the EUV and microwave bands. This study continues our previous work (Sych \\textit{et. al.}, 2009) on the phenomenological relationship between the oscillatory processes in sunspot atmospheres and flaring energy releases, such as flares and spikes. The main emphasis of the present study was put on the analysis of three-minute waves in the transition region and the corona in quite conditions. In agreement with previous studies of this phenomenon (\\textit{e.g.}, Berghmans and Clette, 1999; De Moortel \\textit{et al.}, 2000; De Moortel \\textit{et al.}, 2002a,b; Robbrecht \\textit{et al.}, 2001; King \\textit{et al.}, 2003; De Moortel, 2006), it was found that the power of narrowband three-minute periodic variations of the EUV and microwave emission was localised along the locally open magnetic-field lines (determined by potential-field extrapolation) near their footpoints. The regions of enhanced power of the narrowband periodic variations, called wave traces, highlight the waveguiding magnetised plasma structure. The spatial extent of the wave traces along the magnetic field lines was found to be about 20\\,--\\,30$^{\\prime\\prime}$, again in agreement with previous studies.  The long and short-period oscillations were found to be localised in different regions over the sunspot. A similar feature has been noticed \\textit{e.g.} by Zhugzhda \\textit{et al.} (2000) and Kobanov \\textit{et al.} (2006) at the photospheric and chromospheric levels, where three- and five-minute oscillations were detected at different parts of a sunspot. In the corona, this was established by De Moortel \\textit{et al.} (2002), and Sych and Nakariakov (2008). The three-minute wave traces determined at the transition region and coronal levels are situated at the same locations and have similar shape (taking into account an order-of-magnitude difference in the resolution). It is very likely that we observe the same wave motion of the longitudinal mode, guided along the magnetic field.  The analysis of the temporal evolution of three-minute wave traces revealed its non-stationary nature. The projected phase speed was found to correlate with the wave amplitude. The change of the projected wave speed is likely to be attributed to the change in the magnetic-field geometry, perhaps caused by the waves. In the transition region, the projected phase speed is very high, 180\\,--\\,360 km$^{-1}$, which is much higher than the sound speed at this height. We explain this discrepancy as a projection effect. Indeed, when the longitudinal three-minute waves propagate along the magnetic flux tubes in the corona, they are localised in the transverse direction, and their transverse wavelength is much smaller than the longitudinal wavelength. In other words, the waves are essentially non-planar. This leads to the apparent decrease in the observed projected phase speed with the decrease in the angle between the LOS and the magnetic flux tube, reported in a number of papers (\\textit{e.g.} King \\textit{et al.}, 2003). However, below the corona, the decrease in the sound speed causes the decrease in the longitudinal wavelength. Hence, the three-minute waves become more planar. This leads to the increase in their projected phase speed with the decrease in the angle between the LOS and the direction of their propagation. (We would like to recall that the projected phase speed of a plane wave propagating along the LOS is infinite). Thus, we established that in the transition region the growth of the three-minute amplitude is accompanied by the decrease in the LOS angle to the wave-propagation direction. In other words, the higher amplitude waves are observed to propagate more vertically. A possible interpretation of this phenomenon is the effect of the centrifugal force produced by the field-aligned flows of plasma along the curved magnetic-field lines. This interpretation is qualitatively consistent with the interpretation of the longitudinal compressible coronal waves as slow magnetoacoustic waves (\\textit{e.g.} Nakariakov \\textit{et al.}, 2000; De Moortel, 2006; De Moortel, 2009).  These findings demonstrate the efficiency of the described data analysis system, its ability to bring new interesting scientific results, and its high potential for coronal seismology."
    },
    "1005/1005.1846_arXiv.txt": {
        "abstract": "{We use the SPIRE Fourier-Transform Spectrometer (FTS) on-board the ESA {\\it Herschel} Space Telescope to analyse the submillimetre spectrum of the Ultra-compact H{\\small II} region G29.96-0.02. Spectral lines from species including $^{13}$CO, CO, [CI], and [NII] are detected. A sparse map of the [NII] emission shows at least one other HII region neighbouring the clump containing the UCHII. The FTS spectra are combined with ISO SWS and LWS spectra and fluxes from the literature to present a detailed spectrum of the source spanning three orders of magnitude in wavelength. The quality of the spectrum longwards of 100\\,$\\mu$m allows us to fit a single temperature greybody with temperature $80.3\\pm0.6$\\,K and dust emissivity index $1.73\\pm0.02$, an accuracy rarely obtained with previous instruments. We estimate a mass of 1500~M$_\\odot$ for the clump containing the H{\\small II} region. The clump's bolometeric luminosity of $4\\times10^{6}$\\,L$_\\odot$ is comparable to, or slightly greater than, the known O-star powering the UCH{\\small II} region. } ",
        "introduction": "Ultra-compact H{\\small II} (UCH{\\small II}) regions are small nebulae that surround massive, young stars that are still embedded within a natal cloud \\citep{2002ARA&A..40...27C, 2007prpl.conf..181H}. UCH{\\small II} regions represent an important evolutionary stage in the formation of massive stars. They may correspond to the still embedded accretion phase while the star is already ionising its surroundings. Their study is critical to our understanding of how stars manage to reach high masses despite the high pressure generated by their growing HII regions. UCHII regions are amongst the brightest Galactic objects at submillimetre wavelengths and can be used as a probe of star formation in distant galaxies. Thus, they make excellent candidates with which to probe all phases of massive star formation and with which to test new instruments.  \\object{G29.96-0.02} is an UCH{\\small II} region \\citep{1989ApJS...69..831W,2002A&A...381..571P} located at a distance of $8.9^{+0.6}_{-0.09}$\\,kpc \\citep{2004ApJS..154..553S}. The bright IR source at the centre of G29.96-0.02 is an O5-6 star \\citep{1997ApJ...490L.165W, 2003A&A...405..175M}, with a luminosity of $L_{bol} \\sim 3-4  \\times10^{6}$\\,L$_{\\odot}$ for the \\citet{1997ApJ...490L.165W} temperature limits at the \\citet{2004ApJS..154..553S} distance. However, models indicate that a cluster of young stars must be present in G29.96-0.02 to account for its observed luminosity  \\citep{2003MNRAS.340..799L}.  At the head of the cometary shaped radio emission \\citep{1989ApJS...69..831W} is a small cluster of hot molecular cores ($\\sim$300\\,K) with masses of 3-11\\,M$_{\\odot}$ \\citep{2007A&A...468.1045B}. These arcsec-sized structures are embedded within an arcmin-sized submillimetre clump called \\object{G29.956-0.017SMM} \\citep{2006A&A...453.1003T}. \\citet{1995ApJ...453..308F} suggested that the cometary shape of the UCH{\\small II} region is caused by its expansion westwards into the denser part of the surrounding clump. This clump is part of a larger complex that has been mapped by {\\it Herschel}-SPIRE/PACS as part of the Hi-GAL key programme \\citep[][Beltran et al. in prep.]{2010molinari}.  In this paper we present new submillimetre spectra taken towards \\object{G29.956-0.017SMM} (hereafter G29) with the ESA {\\it Herschel} Space Observatory \\citep{2010herschel} as part of the ``Evolution of Interstellar Dust'' key programme \\citep{2010sag4}. The data were taken with the SPIRE Fourier-Transform Spectrometer (FTS). The SPIRE instrument, its in-orbit performance, and its scientific capabilities are described by \\citet{2010spire}, and the SPIRE astronomical calibration methods and accuracy are outlined by \\citet{2010spirecal}. ",
        "conclusions": "We have presented new SPIRE FTS 190-670$\\mu$m spectra of the submillimetre clump G29.956-0.017SMM which contains the G29.96-0.02 UCH{\\small II} region. The impressive capabilities of the SPIRE FTS have allowed us to simultaneously observe both the dust continuum and prominent spectral lines towards G29. We have conducted basic line-fitting and shown the distribution of [N~{\\sc ii}] emission towards G29. While this preliminary map was not fully sampled it does show at least one other HII region neighbouring G29 and demonstrates the amazing potential of the FTS as a mapping spectrometer. We have reconstructed the SED of G29 using the FTS spectra and archival measurements and have shown that the FTS calibration is broadly consistent with earlier observations. The combined data-set allowed us to fit a precise greybody with $T=80$\\,K and $\\beta=1.73$. Based on a distance of 8.9\\,kpc we estimated the mass of G29 to be approximately $1500$\\,M$_{\\odot}$. The calculated luminosity of the G29 clump is slightly greater than the known O-star at the centre of the UCH{\\small II} region."
    },
    "1005/1005.3302_arXiv.txt": {
        "abstract": "Type I X-ray bursts from low-mass X-ray binaries result from a thermonuclear runaway in the material accreted onto the neutron star. Although typical recurrence times are a few hours, consistent with theoretical ignition model predictions, there are also observations of bursts occurring as promptly as ten minutes or less after the previous event. We present a comprehensive assessment of this phenomenon using a catalog of 3387 bursts observed with the BeppoSAX/WFCs and RXTE/PCA X-ray instruments. This catalog contains 136 bursts with recurrence times of less than one hour, that come in multiples of up to four events, from 15 sources. Short recurrence times are not observed from so-called ultra-compact binaries, indicating that hydrogen burning processes play a crucial role. As far as the neutron star spin frequency is known, these sources all spin fast at over 500~Hz; the rotationally induced mixing may explain burst recurrence times of the order of 10~min. Short recurrence time bursts generally occur at all mass accretion rates where normal bursts are observed, but for individual sources the short recurrence times may be restricted to a smaller interval of accretion rate. The fraction of such bursts is roughly 30\\%. We also report the shortest known recurrence time of 3.8~minutes. ",
        "introduction": "Type I X-ray bursts are thought to result from thermonuclear flashes of hydrogen and/or helium in the envelope of neutron stars (\\citealt{Woosley1976,Maraschi1977,Lamb1978}). This material is accreted through Roche-lobe overflow from a lower-mass companion star (low-mass X-ray binary, LMXB). Current one-dimensional models successfully explain burst features such as the peak flux, the fluence, decay time and recurrence time (e.g., \\citealt{Woosley2004,Heger2007}; see \\citealt{Wallace1981,Fujimoto1981,Fushiki1987ApJ} for earlier work). During the flash, over $90\\%$ of the accreted hydrogen and helium is expected to burn to carbon and heavier elements (e.g., \\citealt{Woosley2004}). For the next flash to occur, a fresh layer of hydrogen/helium must first be accreted. At typical accretion rates of up to approximately $10^{-8}\\,\\mathrm{M_{\\odot}yr^{-1}}$ this takes at least a few hours.  X-ray bursts have been observed since the 1970's (\\citealt{Grindlay1976,1976Belian}) from approximately $90$ sources in our Galaxy, with recurrence times of hours up to days (e.g., \\citealt{Lewin1993,Strohmayer2006}). \\citet{lewin76mnras} reported the detection with \\emph{SAS-3} of three bursts that were separated by only $17$ and $4$~minutes. These bursts originated from a crowded region and, therefore, source confusion cannot be ruled out. In the 1980's similar recurrence times as short as $10$~minutes were observed from both 4U~1608-522 with \\emph{Hakucho} (\\citealt{1608:murakami80pasj}) and from EXO~0748-676 with \\emph{EXOSAT} (\\citealt{0748:gottwald86apj,0748:gottwald87apj}). This rare phenomenon implies that hydrogen and helium is left over somewhere on the star after the initial flash, because the recurrence time is too short to accrete enough fuel for the subsequent burst(s). This is at odds with the current models, that predict an almost complete burning of the available hydrogen and helium on the entire star surface.  \\citet{Boirin2007} analyzed 158 hours of \\emph{XMM-Newton} observations of EXO~0748-676, which revealed short recurrence time bursts in groups of two (doubles) and three (triples). This relatively large burst sample revealed that on average bursts with a short recurrence time ($8$ to $20$~minutes) are less bright and energetic than bursts with `normal' recurrence times (over $2$~hours). The fit of a black body model to the burst spectrum shows a lower peak temperature, while the emitting area is the same. The profiles of short recurrence time bursts seemingly lack the long $50\\,\\mathrm{s}$ to $100\\,\\mathrm{s}$ tail caused by \\textsl{rp}-process burning, which indicates that the burst fuel contains less hydrogen. After a double or triple it takes on average more time before another burst occurs, suggesting a more complete burning of the available fuel.  \\citet{Galloway2008catalog} showed that there are more sources that show this behavior, with bursts occurring in groups of up to four bursts, and with recurrence times as short $6.4$~minutes. The short recurrence times were observed predominantly when the persistent flux is between approximately $2\\%$ and $4\\%$ of the Eddington limited flux. Furthermore, indications were found for the association of short recurrence times and the accretion of hydrogen-rich material. The shortest recurrence time previously reported is $5.4$~minutes (\\citealt{Linares2009ATel}).  Different ideas have been put forward to explain this rare bursting behavior. As most of the models only resolve the neutron star envelope in the radial direction, it is possible that short recurrence time bursts are due to multi-dimensional effects, such as the confinement of accreted material on different parts of the surface, possibly as the result of a magnetic field (e.g., \\citealt{Melatos2005,Lamb2009}). \\citet{Boirin2007}, however, found that the different bursts originate from an emitting area of similar size. Furthermore, the indication of a different fuel composition for the bursts with short recurrence times, argues against any scenario where accreted material of the same composition burns on different parts of the surface.  The idea of a burning layer with an unburned layer on top has been investigated (\\citealt{1636:fujimoto87apj}). After the first layer flashes, the second layer could be mixed down to the depth where a thermonuclear runaway occurs. Mixing may be driven by rotational hydrodynamic instabilities (\\citealt{Fujimoto1988A&A}) or by instabilities due to a rotationally induced magnetic field (\\citealt{Piro2007,Keek2009}). The mixing processes take place on the correct time scale of approximately ten minutes. Although this scenario is able to explain many of the observed aspects of short recurrence time bursts, it has not been reproduced with a multi-zone stellar evolution code that includes a full nuclear burning network. \\citet{Taam1993} created models that exhibit `erratic' bursting behavior, reminiscent of short recurrence time bursts. Later versions of the employed code, however, no longer produce this, most likely because of the inclusion of a more extensive nuclear network (\\citealt{Woosley2004}). As a different explanation for the reignition, \\citet{Boirin2007} suggested a waiting point in the chain of nuclear reactions. There may be a point in the chain where a decay reaction with a half life similar to the short recurrence times stalls nuclear burning before continuing.  We use an improved version of the burst catalog compiled by \\citet{Galloway2008catalog}, that is extended with the X-ray burst observations of the WFCs on-board \\emph{BeppoSAX} (e.g., \\citealt{Cornelisse2003}). This is the largest collection of X-ray bursts used in any study to date. This allows us to study the short recurrence time phenomenon in much more detail. ",
        "conclusions": "We studied thermonuclear bursts with short recurrence times (SWT) using a large catalog of bursts from multiple sources observed with the \\emph{RXTE} PCA and the \\emph{BeppoSAX} WFCs. The short recurrence times are of insufficient duration to accrete the fuel that burns in an SWT burst. Bursts are seen to occur in events of up to four bursts: double, triple and quadruple bursts. We report the shortest recurrence time ever found: $3.8$~minutes. We confirm the result of \\citet{Boirin2007}, that SWT bursts are on average less bright, cooler, and less energetic than LWT bursts. The decay profiles of SWT bursts lack the longer decay component from the \\textsl{rp}-process, suggesting that SWT bursts take place in a hydrogen-depleted layer.  Some sources exhibit short recurrence times at all values of the mass accretion rate where normal bursts occur, while for others SWT bursts are limited to a smaller range of accretion rate. In this range the fraction of bursts with a short recurrence time is consistent with $30\\%$. Two frequent bursters that likely accrete hydrogen-rich matter do not show any SWT bursts.  Only the hydrogen-accreting neutron stars in our catalog exhibit SWT bursts. This suggests that the hydrogen-burning processes are responsible for the incomplete burning of the available hydrogen and helium during bursts. The mechanism for halting the burning is still unknown. It will require further theoretical modeling of hydrogen-accreting neutron star envelopes to resolve this issue.  As far as we know the spin of the sources with SWT bursts, they are all fast rotators. This indicates that rotational mixing can be responsible for ignition of follow-up bursts on a time scale of approximately $10$~minutes (\\citealt{Piro2007,Keek2009}). The number of SWT sources with known spin is small. Measurements of the spin frequency for more sources and further model studies of rotational mixing will better constrain this reignition scenario."
    },
    "1005/1005.2903_arXiv.txt": {
        "abstract": "{Water is a key molecule in the star formation process, but its spatial distribution in star-forming regions is not well known.} {We study the distribution of dust continuum and \\hho\\ and \\thco\\ line emission in DR21, a luminous star-forming region with a powerful outflow and a compact \\ion{H}{ii} region.} {\\textit{Herschel}-HIFI spectra near 1100\\,GHz show narrow \\thco\\ 10$-$9 emission and \\hho\\ $1_{11}-0_{00}$ absorption from the dense core and broad emission from the outflow in both lines. The \\hho\\ line also shows absorption by a foreground cloud known from ground-based observations of low-$J$ CO lines.} {The dust continuum emission is extended over 36$''$ FWHM, while the \\thco\\ and \\hho\\ lines are confined to $\\approx$24$''$ or less. The foreground absorption appears to peak further North than the other components. Radiative transfer models indicate \\new{very low} abundances of $\\sim$\\pow{2}{-10} for \\hho\\ and $\\sim$\\pow{8}{-7} for \\thco\\ in the dense core, \\new{and higher \\hho\\ abundances of $\\sim$\\pow{4}{-9} in the foreground cloud and $\\sim$\\pow{7}{-7} in the outflow}.} {The high \\hho\\ abundance in the warm outflow is probably due to the evaporation of water-rich icy grain mantles, \\new{while the \\hho\\ abundance is kept down by freeze-out in the dense core and by photodissociation in the foreground cloud}.} ",
        "introduction": "\\label{s:intro}  The water molecule is a key species throughout the formation of stars and planets. In the gas phase, it acts as a coolant of collapsing interstellar clouds; in the solid state, it acts as glue for dust grains in protoplanetary disks to make planetesimals; and as a liquid, it acts as transporter bringing molecules together on planetary surfaces, a key step towards biogenic activity. The first role is especially important for high-mass star formation which depends on the balance between the collapse of a massive gas cloud and its fragmentation \\citep{review}.  Interstellar \\hho\\ is well known from ground-based observations of the 22\\,GHz maser line. % Previous space-based submm and far-IR observations have measured \\hho\\ abundances ranging from 10$^{-8}$ in cold gas to 10$^{-4}$ in warm gas (ISO: \\citealt{iso}; SWAS: \\citealt{swas}; Odin: \\citealt{odin}) but did not have sufficient angular resolution to determine the spatial distribution of \\hho. In contrast, space-based mid-IR and ground-based mm-wave observations have high angular resolution but only probe the small fraction of the gas at high temperatures (\\citealt{iram}; \\citealt{spitzer}).  This paper presents observations of an \\hho\\ ground state line at $>$3$\\times$ higher angular resolution than previously possible for such lines. Through radiative transfer models, we compare the abundance distribution of \\hho\\ with that of \\thco\\ and dust. The source DR21 (Main) is a high-mass protostellar object ($L$=45,000\\,\\lsol) located in the Cygnus~X region at $d$=1.7\\,kpc \\citep{schneider}, about 3$'$ South of the well-known DR21(OH) object (also known as W75S). Maps of the 1.2\\,mm dust emission show a dense core with a mass of 600--1000\\,\\msol\\ and a size of 0.19$\\times$0.14\\,pc FWHM, surrounded by an extended envelope with mass 4750\\,\\msol\\ and size 0.3~pc \\citep{motte}. Gas densities of 10$^5$--10$^6$\\,\\ccm\\ are derived from both the mm-wave continuum and HCN and HCO$^+$ line emission \\citep{kirby}. Signs of active high-mass star formation are the bright mid-IR emission (272\\,Jy at 21\\,$\\mu$m), the presence of an \\hho\\ 22\\,GHz maser (see catalog of \\citealt{braz}) and emission from ionized gas extending over 20--30$''$ \\citep{pjotr}. Together with the powerful molecular outflow \\citep{garden} these signs indicate that the source is relatively evolved within the embedded phase of high-mass star formation, beyond the `ultracompact \\ion{H}{ii} region' phase.  \\begin{figure*}[tb] \\centering \\includegraphics[width=6cm,angle=0]{14515f1a.ps} \\includegraphics[width=6cm,angle=0]{14515f1b.ps} \\caption{Spectra of \\thco\\ $10-9$ (left) and \\hho\\ $1_{11}-0_{00}$ (right) lines toward DR21, taken with the WBS backend.} \\label{f:wbs} \\end{figure*} ",
        "conclusions": "\\label{s:disc}   To estimate the \\hho\\ and \\thco\\ abundances from our data, we have run spherical radiative transfer models following \\citet{marseille}. First, the dust continuum emission was modeled with the MC3D program \\citep{mc3d} with the source size and luminosity kept fixed at the values in \\S\\ref{s:intro}. The continuum data are consistent with a power-law density profile $n=n_0 (r/r_0)^{-\\alpha}$ with the index $\\alpha$=1.5 as expected for evolved protostellar envelopes \\citep{massive}. Derived temperatures range from 117\\,K at the adopted inner radius of 0.01~pc to 23\\,K at the outer radius of 0.3~pc; densities drop from \\pow{3}{7}\\,\\ccm\\ to \\pow{2}{5}\\,\\ccm. This temperature and density profile was adopted for the line radiative transfer with the RATRAN program \\citep{ratran}. The abundance of \\hho\\ was varied between 10$^{-10}$ and 10$^{-7}$ and the \\thco\\ abundance between 10$^{-7}$ and 10$^{-6}$, both independent of radius.  \\begin{figure}[tb] \\centering \\includegraphics[width=3.5cm,angle=-90]{14515f4a.eps} \\includegraphics[width=4.5cm,angle=-90]{14515f4b.eps} \\caption{Data from Fig.~\\ref{f:hrs} with radiative transfer models superposed. } \\label{f:ratran} \\end{figure}   The red lines in Figure~\\ref{f:ratran} show the results of our best-fit model. The fit to the line profiles at the central few positions is good if two components are added to the dense core model: one representing the outflow and one for the foreground cloud. This simple addition procedure is valid because the foreground cloud is transparent and the outflow is seen at an angle, so that neither component blocks the view of the dense core. The modeled optical depths are 2.8 for the \\hho\\ line and 0.77 for \\thco, and the derived abundances are uncertain to $\\approx$35\\% which is the quadratic sum of 30\\% from the line intensity and 20\\% from the core mass estimate. The model underproduces the observed line strengths at large position offsets, which suggests that the density profile flattens out at large radii, as also indicated by the observed continuum brightness at the offset positions.  Deriving the \\thco\\ and \\hho\\ abundances in the outflow and the foreground using radiative transfer models is not possible because the masses and \\hh\\ column densities of these components are unknown. Instead we have used RADEX \\citep{radex} to estimate their \\thco\\ and \\hho\\ column densities, \\new{and use a $^{12}$C/$^{13}$C ratio of 60 \\citep{isotope} and a CO abundance of \\pow{2}{-4} \\citep{lacy} to convert $N$(\\thco) to $N$(\\hh) and estimate $x$(\\hho) as $N$(\\hho)/$N$(\\hh)}. For the foreground cloud, we adopt \\tkin = 10\\,K and \\nhh = 10$^4$\\,\\ccm, and for the outflow, we adopt \\tkin = 200\\,K and \\nhh = \\pow{3}{4}\\,\\ccm, \\new{which assumptions introduce a factor of $\\sim$2 uncertainty each}. To estimate $N$(\\thco) for the foreground cloud, we use the $J$=1-0 observations by \\citet{jakob}.  Table~\\ref{t:abs} summarizes our derived column densities and abundances of \\thco\\ and \\hho\\ in the various physical components of the DR21 region. \\new{Our \\hho\\ abundance in the dense core is $\\sim$100$\\times$ lower than previous determinations (\\S\\ref{s:intro}) but should be regarded as a lower limit.} At the low temperatures and high densities in the core, most \\hho\\ is likely frozen on grains, and the observed line may arise in a small region with a high \\hho\\ abundance. The derived \\thco\\ abundance for the core is $\\sim$4$\\times$ lower than expected for the above values of the CO isotopic ratios and abundance, which suggests that even some CO is frozen out in the outer parts of the core. The density of the foreground cloud is too low for significant freeze-out, but with $A_V$ $\\approx$1.2\\,mag, photodissociation is rapid for \\hho\\ but not for \\thco.  The high \\hho\\ abundance for the outflow is likely related to its temperature of $\\sim$200\\,K, which is high enough to have \\hho\\ released from the dust grains by thermal evaporation, or possibly by shocks \\citep{melnick}. Further enhancement may be expected in even warmer gas ($\\gtsim$250\\,K) when neutral-neutral reactions drive most gas-phase oxygen into \\hho, but such gas is not probed by our data. Future HIFI observations of high-excitation \\hho\\ lines towards protostars of all masses will however very likely reveal this effect.  \\begin{table} \\caption{Abundances and column densities of \\thco\\ and \\hho.} \\label{t:abs} \\begin{tabular}{ccccc} \\hline \\hline \\noalign{\\smallskip} Component & $N$(\\thco)$^a$  & $N$(p-\\hho)     & $N$(\\hh) & $x$(p-\\hho)\\\\ & 10$^{16}$\\,\\scm & 10$^{12}$\\,\\scm & 10$^{21}$\\,\\scm & \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Dense core & \\pow{7.8}{-7} & ...              & ... & \\pow{1.6}{-10} \\\\ Outflow    & 5             & 10,000           & 15  & \\pow{7}{-7} \\\\ Foreground & 0.7           & 4                & 2.1 & \\pow{4}{-9} \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{5}{l}{$^a$: Abundance for the dense core} \\\\ \\end{tabular} \\end{table}"
    },
    "1005/1005.2768_arXiv.txt": {
        "abstract": "We study a simple class of time-dependent rotating Ricci-flat cylindrically symmetric spacetime manifolds whose geodesics admit gravitomagnetic jets. The helical paths of free test particles in these jets up and down parallel to the rotation axis are analogous to those of charged particles in a magnetic field. The jets are attractors. The jet speed asymptotically approaches the speed of light. In effect, such source-free spacetime regions act as ``gravitomagnetic accelerators.'' ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.0840_arXiv.txt": {
        "abstract": "{ We derived the absolute proper motion (PM) of the old, solar-metallicity Galactic open cluster M~67 using observations collected with CFHT (1997) and with LBT (2007).  About 50 galaxies with relatively sharp nuclei allow us to determine the absolute PM of the cluster.  We find $(\\mu_\\alpha\\cos\\delta,\\,\\mu_\\delta)_{\\rm J2000.0}$ $=$ $(-9.6\\pm1.1,\\,-3.7\\pm0.8)$ \\masyr.  By adopting a line-of-sight velocity of $33.78 \\pm 0.18$ \\kms, and assuming a distance of 815$\\pm$50 pc, we explore the influence of the Galactic potential, with and without the bar and/or spiral arms, on the galactic orbit of the cluster. } ",
        "introduction": "\\label{sec:1}   The solar-metallicity Galactic open cluster M~67 (NGC~2682) is among the most-studied Galactic open clusters. Still, its absolute proper motion (PM) remains poorly constrained.  We applied for the first time on ground-based multi-epoch CCD wide-field images the PM techniques developed in Anderson et al.\\ (2006, Paper~I), Yadav et al.\\ (2008, Paper~II), Bellini et al.\\ (2009, Paper~III) to define an absolute reference frame using background faint galaxies.  In the next section we describe the data set and the measurements, and a final section is dedicated to the study of the orbit of M 67 within the Galaxy under different assumptions for the Galactic potential. ",
        "conclusions": ""
    },
    "1005/1005.2845_arXiv.txt": {
        "abstract": "{ The discovery of a historical bug in the s-post-process AGB code obtained so far by the Torino group forced us to reconsider the role of primary $^{16}$O in the $^{13}$C-pocket, produced by the $^{13}$C($\\alpha$, n)$^{16}$O reaction, as important neutron poison for the build up of the s-elements at Halo metallicities. The effect is noticeable only for the highest $^{13}$C-pocket efficiencies (cases ST*2 and ST). For Galactic disc metallicities, the bug effect is negligible. A comparative analysis of the neutron poison effect of other primary isotopes ($^{12}$C, $^{22}$Ne and its progenies) is presented. The effect of proton captures, by  $^{14}$N(n, p)$^{14}$C, boosts a primary production of Fluorine in Halo AGB stars, with [F/Fe] comparable to [C/Fe], without affecting the s-elements production. ",
        "introduction": " ",
        "conclusions": "For Halo metallicities, the $^{16}$O bug we had in the $^{13}$C-pocket, which consisted in discarding neutron capture on the $^{16}$O isotopes released by $^{13}$C($\\alpha$, n)$^{16}$O, affects [ls/Fe], [hs/Fe], [Pb/Fe] in an important way for the highest choices of the $^{13}$C-pocket efficiency (cases ST $\\times$ 2 and ST). Instead, the s-process indexes [hs/ls], [Pb/hs] are only marginally affected. For Galactic disc AGBs, the bug effect is negligible for all $^{13}$C-pocket choices. Introduction of proton captures in the code does not affect the s-process abundances. However, it boosts  the prediction of Fluorine in very metal-poor s-enhanced stars. A primary production of F, comparable with the C enhancement, is predicted in AGB stars of low metallicity. The only spectroscopic observation of Fluorine in a CEMP-s star, is the giant HE 1305+0132 by Schuler et al. (2007, 2008), and is quite in agreement with our F prediction. Higher resolution spectra are  however necessary for a definite confirmation."
    },
    "1005/1005.2060_arXiv.txt": {
        "abstract": "We present first observations of a dome-shaped large-scale EUV coronal wave, recorded by the EUVI instrument onboard STEREO-B on January 17, 2010. The main arguments that the observed structure is the wave dome (and not the CME) are: a) the spherical form and sharpness of the dome's outer edge and the erupting CME loops observed inside the dome;  b) the low-coronal wave signatures above the limb perfectly connecting to the on-disk signatures of the wave; c) the lateral extent of the expanding dome which is much larger than that of the coronal dimming; d) the associated high-frequency type II burst indicating shock formation low in the corona. The velocity of the upward expansion of the wave dome ($v \\sim 650$~km~s$^{-1}$) is larger than that of the lateral expansion of the wave ($v \\sim 280$~km~s$^{-1}$), indicating that the upward dome expansion is driven all the time, and thus depends on the CME speed, whereas in the lateral direction it is freely propagating after the CME lateral expansion stops. We also examine the evolution of the perturbation characteristics: First the perturbation profile steepens and the amplitude increases. Thereafter, the amplitude decreases with r$^{-2.5 \\pm 0.3}$, the width broadens, and the integral below the perturbation remains constant. Our findings are consistent with the spherical expansion and decay of a weakly shocked fast-mode MHD wave. ",
        "introduction": "Large-scale propagating disturbances in the solar atmosphere occuring in association with flares and coronal mass ejections (CMEs) have been first observed in chromospheric filtergrams \\citep{moreton60,athay61}. These ``Moreton waves'' propagate with typical velocities in the order of 1000~km~s$^{-1}$, and have been interpreted as the ground-track of a dome-shaped MHD wave front propagating through the solar corona, which compresses and pushes the chromospheric plasma downward when sweeping over it \\citep{uchida68}. % The Extreme-ultraviolet Imaging Telescope \\citep[EIT;][]{delaboudiniere95} onboard SOHO for the first time imaged such wave-like disturbances in coronal emission lines \\citep{moses97,thompson98}.  EIT waves have been initially interpreted as the coronal counterparts of the chromospheric Moreton waves as predicted in Uchida's fast-mode coronal MHD wave model \\citep{thompson99}.  % However, differences in morphology and propagation velocities of Moreton waves and EIT waves, which lie mostly in the range 200--400 km~s$^{-1}$ \\citep{klassen00,thompson09} led to severe doubts of this interpretation and alternative models were put forward. Some of them question if the phenomenon is a wave at all, and instead suggest that it is a signature of the large-scale coronal restructuring due to the erupting CME causing plasma compression or localized energy release \\citep[e.g.][]{delanee99,chen02,attrill07}. For detailed discussions of the different models, we refer to the recent reviews by \\cite{warmuth07}, \\cite{vrsnak08}, and \\cite{wills10}.   There seems to be some consensus that Moreton waves are indeed shock waves, as is suggested by their high propagation velocities and perturbation characteristics. It was shown that the amplitude and velocity of Moreton waves decrease, and the width of the wave pulse broadens as it propagates, consistent with a large-amplitude wave or freely propagating shock wave that formed by steepening of a simple wave \\citep[e.g.][]{warmuth04b}. However, the situation is unclear for large-scale waves observed in the corona. There is rather qualitative insight that the wave fronts not only become more diffuse but also broaden as they propagate \\citep{thompson99,klassen00,podladchikova05}, which may be due to energy flux conservation or dispersion of the wave  \\citep[e.g.][]{wills10}. On a statistical basis of different EIT waves, \\cite{warmuth10} concludes that for larger distances, the perturbation amplitudes tend to become smaller and the width larger. However, a case study of the wave evolution in high-cadence TRACE images by \\cite{wills06} revealed that the width of the wave pulse remained constant during the propagation.  Due to the low EIT cadence of $\\sim$12--15~min it was not possible to detect fast EIT waves and to study the evolution of the wave pulse characteristics, which is important in constraining the physical processes. The Extreme Ultraviolet Imager \\citep[EUVI;][]{howard08} instruments onboard the twin spacecraft of the Solar-Terrestrial Relations Observatory \\citep[STEREO;][]{kaiser08} overcome these limitations and offer several advantages for the study of large-scale coronal waves, in particular due to their high cadence, large field-of-view (FoV), high sensitivity, and the simultaneous observations from two vantage points.  Since the launch of the STEREO satellites in October 2006, a variety of case studies of large-scale coronal waves have been carried out using EUVI data \\citep{long08,veronig08,gopalswamy09,patsourakos09,patsourakos09b,attrill09,cohen09,kienreich09,ma09,zhukov09,dai10}. The typical propagation velocities that were derived from EUVI waves (that all occurred during solar minimum conditions and were not accompanied by Moreton waves) are in the range 200--350~km~s$^{-1}$. For the kinematical evolution, the results so far are rather inconclusive: some of the waves studied with EUVI showed evidence for deceleration during their propagation, consistent with the decay and deceleration of a large amplitude wave to the fast-mode speed of the ambient corona \\citep[e.g.][]{long08,veronig08}, whereas others showed constant velocity \\citep[e.g.][]{ma09,kienreich09}. It is worth noting that many of the EUVI wave studies revealed a close association of the wave and the erupting CME and its expanding flanks \\citep[e.g.][]{veronig08,attrill09,kienreich09,patsourakos09b}, whereas the associated flares were all very weak. Stereoscopic EUVI studies revealed that the wave signal is typically confined to about 1--2 coronal scale heights above the chromosphere \\citep{patsourakos09,patsourakos09b,kienreich09}.  In this letter, we present the first observations of the full dome of the wave observed in EUV, which is consistent with the three-dimensional expansion of a coronal shock front. We note that part of a wave dome had been observed in soft X-rays by Yohkoh/SXT for a fast coronal wave that occurred in association with a Moreton wave \\citep{narukage04}. We also study the evolution of the wave pulse characteristics (velocity, amplitude, width) in high-cadence EUV imaging over a propagation distance of more than $1\\,R_\\odot$, and discuss the implications for the physical processes involved. ",
        "conclusions": "There are four main arguments supporting the conclusion that the observed EUV structure is the wave dome and not the CME body: a) The dome appears spherically three-dimensional with a sharp outer edge; inside the dome erupting CME loops are observed. b) The wave signatures observed in the low corona above the limb perfectly connect to the wave signatures observed against the solar disk. c) The lateral extent of the expanding dome is much larger than that of the coronal dimming. d) The event was associated with a high-frequency type II burst indicating shock formation low in the corona, which is consistent with the observed EUV dome height.  The velocity derived for the upward expansion of the wave dome  ($v \\sim 650$~km~s$^{-1}$) is larger than that of the lateral expansion of the wave observed on-disk ($v \\sim 280$~km~s$^{-1}$). There are at least two alternative explanations for this. In the freely propagating phase, the upward-lateral velocity difference can be due to differences of the fast magnetosonic velocities of the active region (AR) and the ambient quiet corona. The tip of the wave dome lies above the AR, and the magnetosonic speed is much higher above ARs than at low heights in its surroundings \\cite[e.g.][]{warmuth05b}. An alternative explanation is that the upward dome expansion is driven all the time, and thus depends on the CME speed, whereas in the lateral direction the wave is freely propagating as soon as the lateral expansion of the CME flanks has stopped, and its velocity is determined by the characteristic speed of the medium. The second interpretation is supported by the evolution of the dimming region in the low corona, which is fast expanding up to about 04:01--04:06~UT. Thereafter, the coronal dimming still gets darker until about 05:00~UT, indicating ongoing mass depletion, but with only little (or no) further expansion after 04:06~UT.  The high-cadence EUV observations together with the distinct wave signal observed in the low corona against the solar disk allowed us to study in detail the evolution of the perturbation characteristics during the coronal wave propagation. We find that the amplitude of the perturbation first increases and the perturbation profile steepens within the first 5~min of the event (i.e.\\ until 04:01~UT), thereafter the amplitude~$A$ decreases, following a power-law of the form $A \\propto r^{-2.5 \\pm 0.3}$. The width of the perturbation profile broadens during its evolution by a factor of 3--4. This broadening is observed for both the full width as well as for the width at half maximum. We stress that we only measured the frontal width of the wave pulse due to the uncertainties in the trailing part (mostly due to stationary brightenings induced by the wave passage). Since the perturbation profile is not necessarily symmetric, the evolution of the total width of the wave pulse may be different, but we do not expect that it would alter the main outcome of broadening, which is a quite pronounced effect. The integral below the frontal part of the perturbation profile basically remains constant. These characteristics of the wave pulse evolution are all consistent with the spherical expansion and decay of a large-amplitude fast-mode MHD wave.  The velocity of the lateral wave expansion ($v \\sim 280$~km/s) lies in the range of the fast magnetosonic speed in the quiet solar corona during solar minimum condition \\citep[see discussion in][]{kienreich09}. From the observed EUV wave intensities at the highest amplitude reached at $\\sim$04:01~UT, we obtained a rough estimate of the peak density jump at the wave front, $n/n_0 \\sim 1.2$. This corresponds to a perpendicular fast magnetosonic wave number of $M_{fms} \\sim 1.15$ \\citep{priest82}. Thereafter, as the wave decays to $n/n_0=1.04$ (Fig.~\\ref{fig5}; last profile at 04:36~UT), it approaches a linear regime since $M_{fms} \\sim 1.03$. These numbers indicate the evolution of a weak shock, and may explain why we do not observe a significant deceleration of the free-propagation of the lateral wave expansion in association with the amplitude decay."
    },
    "1005/1005.5476_arXiv.txt": {
        "abstract": "{  By combining {\\it Herschel}-SPIRE data with archival {\\it Spitzer}, \\hi, and CO maps, we investigate the spatial distribution of gas and dust in the two famous grand-design spirals M\\,99 and M\\,100 in the Virgo cluster. Thanks to the unique resolution and sensitivity of the {\\it Herschel}-SPIRE photometer, we are for the first time able to measure the distribution and extent of cool, submillimetre (submm)-emitting dust inside and beyond the optical radius. We compare this with the radial variation in both the gas mass and the metallicity. Although we adopt a model-independent, phenomenological approach, our analysis provides important insights. We find the dust extending to at least the optical radius of the galaxy and showing breaks in its radial profiles at similar positions as the stellar distribution. The colour indices $f350/f500$ and $f250/f350$ decrease radially consistent with the temperature decreasing with radius. We also find evidence of an increasing gas to dust ratio with radius in the outer regions of both galaxies. } ",
        "introduction": "\\begin{figure*} \\hspace*{1cm} SDSS u-z \\hspace{1.5cm} 70 $\\mu$m \\hspace{2.1cm} 250 $\\mu$m \\hspace{1.9cm} 350 $\\mu$m \\hspace{2.0cm} 500$\\mu$m \\hspace{2.3cm} \\hi \\\\[-0.15cm] \\includegraphics[height=18.5cm,angle=270]{14554fg1.ps} \\vspace*{-0.0cm} \\caption{The matched SDSS, MIPS 70 $\\mu$m, SPIRE 250, 350, 500 $\\mu$m, and \\hi-maps of M\\,99 {\\it (upper row)} and M\\,100 {\\it (lower row)}. The {\\it green contours} are SPIRE 250 at levels about 0.03($\\approx\\!5\\sigma$), 0.10, 0.34, 1.15, 3.9 Jy beam$^{-1}$. The 350 and 500 $\\mu$m are displayed at a high contrast to show the extent of the dust emission and the scan-artefact free smooth background dominated only by confusion noise. The {\\it white circle} on the \\hi-maps marks the optical radius defined by $D25$. The {\\it blue circles} on each map indicate the FWHM. Note, the SDSS jpg is scaled but not WCS matched.} \\label{maps} \\vspace*{-0.3cm} \\end{figure*}  The study of the gas and dust distribution in galaxies is essential to understanding their formation and evolution. The rate at which gas is accreted and converted into stars regulates not only the star formation history of galaxies but also their chemical evolution. Dust is supposed to play a key role in this process. Dust grains are the main coolant in star-forming galaxies, shielding the gas from the UV radiation and representing the site at which \\hi\\ is converted into \\hii, and then collapses into stars \\citep[see e.g., reviews by][]{Calzetti2001,Draine2003}.  To understand the dust, we need to map the cold component, which does not dominate the energy but dominates the mass. However, using previously existing facilities, our knowledge of the interplay between gas and dust and their radial distribution have remained highly uncertain, being based only on $\\lambda\\!<\\!  160\\mu$m space observations \\citep[e.g.~][]{Munoz2009a, Bendo2010a}, and challenging ground-based submm observations that were of optimal quality only at significantly longer wavelength. For many galaxies, only integrated quantities have been derived because of the poor resolution of previous satellites. Resolved submm studies from the ground remain limited to the very nearby universe and large surveys of more distant galaxies are unfeasible.  The SPIRE instrument \\citep{Griffin2010} on-board {\\sl Herschel} \\citep{Pilbratt2010} now bridges this gap observing in the range of 250$\\mu$m-500$\\mu$m. With the benefit of the stable conditions of a space observatory, it is much more sensitive to the cold component and provides excellent maps at high resolution, so is ideal for large surveys of many galaxies. With {\\sl Herschel}, we are now able to tackle the problem of the interplay between gas and dust, and combined with the number of recent high resolution surveys tracing the gas mass of local galaxies \\citep[e.g.~][]{Chung2009,Kuno2007}, we can study the distribution of gas, metals, and dust on a kpc scale for hundreds of galaxies.  Here, we discuss first results for two famous grand-design spirals M\\,99 and M\\,100 (see \\fig\\ref{maps}) in the Virgo cluster. We explore the distribution of the cool dust traced with {\\sl Herschel} by inspecting their radial profiles from mid-infrared to submm wavelengths. We then attempt to correlate this with the observed gas and metallicity distributions and search for temperature variations. We use the {\\sl Herschel}-SPIRE maps taken during {\\sl Herschel}'s Science Demonstration Phase. The two galaxies are part of the {\\sl Herschel} Reference Survey \\citep{Boselli2010a}. This guaranteed time key project will provide maps for a statistically-complete sample of 323 nearby galaxies in all three SPIRE bands. In the RC3 \\citep{rc3}, the classification and optical radius $R25$ of M\\,99 and M\\,100 is given as SA(s)c, 2.69\\arcmin ($\\approx 12.9$kpc), and SAB(s)bc, 3.71\\arcmin($\\approx 17.8$kpc), respectively. We assume a distance for the Virgo Cluster of 16.5 Mpc \\citep{Mei2007}. ",
        "conclusions": ""
    },
    "1005/1005.0779_arXiv.txt": {
        "abstract": "We present the gravitational wave signatures for a suite of axisymmetric core collapse supernova models with progenitors masses between 12 and 25 M$_\\odot$. These models are distinguished by the fact they explode and contain essential physics (in particular, multi-frequency neutrino transport and general relativity) needed for a more realistic description. Thus, we are able to compute complete waveforms (i.e., through explosion) based on non-parameterized, first-principles models. This is essential if the waveform amplitudes and time scales are to be computed more precisely. Fourier decomposition shows that the gravitational wave signals we predict should be observable by AdvLIGO across the range of progenitors considered here. The fundamental limitation of these models is in their imposition of axisymmetry. Further progress will require counterpart three-dimensional models. ",
        "introduction": "Core collapse supernovae are among the sources that will produce gravitational waves (GWs) detectable by GW observatories around the globe. In particular, a Galactic supernova event will likely produce a signal well within Advanced LIGO's \\cite{AdvLIGO} bandpass across a broad range of frequencies. Gravitational waves from supernovae will arise from a variety of phenomena given their multidimensional, multi-physics character. These phenomena include fluid instabilities in the proto-neutron star, neutrino-driven convection beneath the supernova shock wave, the standing accretion shock instability (SASI), deceleration at an aspherical shock, and aspherical neutrino emission. Obviously, two-, and ultimately three-, dimensional models will be required to capture the GW emission from such phenomena. Moreover, the explosion dynamics, and ultimately the computation of the GW emissions as well, will require sufficient realism in the treatment of core collapse supernova multi-physics. For a comprehensive survey of the field see the recent review by Ott {\\cite{Ott2009}} and references therein. ",
        "conclusions": "We present gravitational waveforms computed in the context of 2D core collapse supernova simulations performed with the CHIMERA code for non-rotating 12, 15, and 25 $M_\\odot$ progenitors. We calculate the contribution to the signals produced by both baryonic matter motion and anisotropic neutrino emission up to 530 ms after bounce for all three progenitors. Given the development of non-parameterized explosions in our models, we are able to compute the waveforms through explosion and to determine more precisely the pre-explosion amplitudes and timescales. Given our use of tracer particles, we are able to decompose the GW signatures and determine which phenomena contribute to specific components of the waveforms. This allowed us to identify an additional source for the prompt signal (in the past solely attributed to prompt convection) the deflection of infalling matter through the shock.  Our waveforms exhibit a characteristic signature. Namely, the signal develops in four stages. There is 1) a relatively short and weak {\\bf prompt signal}, 2) a {\\bf quiescent stage}, 3) a {\\bf strong signal} where most of the GW energy is emitted and, 4) a slowly increasing {\\bf tail}. We predict signatures with sufficient strength to be readily observable by Advanced LIGO for a Galactic event, and the peak in the observable spectrum stems from the accretion downflows driven by the SASI.  The results presented here are preliminary: a new set of 2D simulations performed with an enhanced version of our CHIMERA code is currently ongoing \\cite{Bruenn10}. However, while a number of approximations are made in the CHIMERA code, the more important limitation in these models is their restriction to axisymmetry. Three-dimensional models are required. We anticipate that the greatest change to our gravitational waveform predictions in moving to 3D will be in the phase-4 tail. Prolate explosions are often seen in axisymmetric simulations, where artificial boundary conditions must be imposed that prevent the turnover of material along the symmetry axis. With axisymmetry removed, we expect a significant change in the evolution of the explosion tail: in its magnitude, and perhaps even in its sign. And no doubt there will be quantitative changes to the amplitudes and timescales associated with earlier phases, particularly in the strong signal arising from the SASI motions. In 3D, the SASI will likely be dominated by spiraling flows \\cite{Blondin07}, fundamentally different than the sloshing modes that dominate in the axisymmetric case. This will in turn alter the waveforms in the final pre-explosion phase. This has already been demonstrated by Kotake \\etal \\cite{Kotake09} in 3D parameterized studies. 3D simulations with all of the physics documented here are ongoing, and we look forward to reporting on their GW signatures in the near future.  {\\bf Acknowledgements:} The authors would like to acknowledge the computational resources provided at the Leadership Computing Facility in the National Center for Computational Sciences at ORNL (INCITE Program) and at TACC (TG-MCA08X010). PM acknowledges partial support from NSF-PHYS-0855315, and PM and SWB acknowledge partial support from an NSF-OCI-0749204 award. AM, OEBM, PM, SWB, and WRH acknowledge partial support from a NASA ATFP award (07-ATFP07-0011). AM and WRH acknowledge support from the Office of Nuclear Physics, U.S. Department of Energy, and AM and OEBM acknowledge support from the Office of Advanced Scientific Computing Research, U.S. Department of Energy."
    },
    "1005/1005.2917_arXiv.txt": {
        "abstract": "The stalling radius of a merging massive binary black hole (BBH) is expected to be below 0\\farcs1 even in nearby galaxies \\citep{yu02}, and thus BBHs are not expected to be spatially resolved in the near future. However, as we show below, a BBH may be detectable through the significantly anisotropic stellar velocity distribution it produces on scales 5--10 times larger than the binary separation. We calculate the velocity distribution of stable orbits near a BBH by solving the restricted three body problem for a BBH embedded in a bulge potential. We present high resolution maps of the projected velocity distribution moments, based on snapshots of $\\sim 10^8$ stable orbits. The kinematic signature of a BBH in the average velocity maps is a counter rotating torus of stars outside the BBH Hill spheres. The velocity dispersion maps reveal a dip in the inner region, and an excess of 20--40\\% further out, compared to a single BH of the same total mass. More pronounced signatures are seen in the third and fourth Gauss--Hermite velocity moments maps. The detection of these signatures may indicate the presence of a BBH currently, or at some earlier time, which depends on the rate of velocity phase space mixing following the BBH merger. ",
        "introduction": "The discovery that most galaxies harbour a massive black hole (BH) at their core \\citep{1998AJ....115.2285M}, and the commonly accepted interpretation of cosmological structure formation simulations, that galaxies grow by mergers (e.g. \\citealt{1993MNRAS.264..201K}, cf. \\citealt{2006MNRAS.368....2D} that growth is mostly by gas accretion) implies that massive binary black holes should commonly form at the core of galaxies \\citep{2003ApJ...582..559V}. Unlike the previous view, that the formation of such a binary results from the rare process of merging during the life of the rare objects, quasars \\citep*{b80}. Dynamical friction on the background stars leads to inspiral of the binary on a dynamical time-scale, until the binary becomes hard, which occurs at a binary separation close to \\begin{equation} a_{\\rm h}\\equiv\\frac{q}{1+q}\\frac{GM_\\bullet}{4\\sigma^2},\\label{eq:ah} \\end{equation} where $G$ is the gravitational constant, $M_\\bullet$ is the mass of the primary BH, $q$ is the mass ratio of the BHs ($q\\leq1$), and $\\sigma$ is the bulge 1D velocity dispersion. The inspiral rate slows drastically when $a \\lesssim a_{\\rm h}$ (e.g. \\citealt{2006ApJ...648..976M}), and further inspiral is set by the rate at which stars diffuse in phase space into the loss cone \\citep{1976MNRAS.176..633F}\\footnote{More properly termed the loss cylinder, as pointed out by \\citet{1978ApJ...226.1087C}. See also \\S\\ref{sec:stability-maps} below.}. The minimum phase space diffusion rate is set by the two body relaxation rate. This mechanism may be fast enough in the lowest luminosity bulges, where the stellar cores are the densest, and may lead to merger through gravitational wave emission in less than the Hubble time (but see recent suggestions that the core density in the Milky Way is smaller than originally thought, \\citealt{2010arXiv1001.5435M} and references therein) . In more luminous bulges, the binary inspiral is expected to stall (e.g. \\citealt{2004ApJ...602...93M}; \\citealt{2006RPPh...69.2513M}, and references therein), leading to the  final parsec problem \\citep{b80}. The largest plausible stalling radius in the nearest galaxies is just below $0\\farcs1$ \\citep{yu02}, and is thus generally unresolved. The stalling may be overcome by diffusion of orbits into the loss cone induced by tangential forces. Either due to bulge triaxiality or bar like structure, or by massive perturbers in the form of molecular clouds or a third BH due to another merger (e.g. \\citealt{2004ApJ...606..788M}; \\citealt{2006ApJ...642L..21B}; \\citealt{2007MNRAS.377..957H}; \\citealt{tal07}; \\citealt{2008ApJ...677..146P}). Alternatively, inspiral may be induced by angular momentum loss of the BBH to circumbinary gas (e.g. \\citealt{2005ApJ...634..921A}; \\citealt{2007Sci...316.1874M}; \\citealt{2009MNRAS.393.1423C}). These processes may lead to a merger on time-scales well below the Hubble time.  The formation process of the BBH ejects stars from the bulge. This occurs on scales significantly larger than $a_{\\rm h}$, and it flattens the stellar density profile. This BBH ``scouring'' may be responsible for the formation of the core present in luminous ellipticals (\\citealt{mm01}; \\citealt{mm05}). Recent high quality observations by \\citet{2009ApJ...691L.142K} indicate a remarkably tight correlation between the deduced mass deficiency in the core region, and the black hole mass. This correlation is likely a signature of the cumulative scouring effects of merging BHs during the BH growth. Are there any other signatures induced by the merger?  If the binary stalls, then a significant fraction of galaxies may harbour a compact BBH (e.g. \\citealt{2003ApJ...582..559V}), which cannot currently be directly resolved \\citep{yu02}. The binary affects the stability of stellar orbits on scales up to several times larger than the binary separation. The presence of a binary may thus be inferred through its signature on the background stellar distribution and kinematics (e.g. \\citealt{2005CQGra..22S.347C}; Kandrup et al. 2003). The purpose of this paper is to determine this signature and provide a tool which can be used to detect the effect of a binary black hole on scales 5--10 times larger than $a_{\\rm h}$, a scale which may be resolved in nearby galaxies with current angular resolutions.  We first explore the stability of orbits near a BBH by calculating stability maps. These maps present the time it takes for a test particle to become unbound, as a function of its position in velocity space, at a given launching radius. A similar concept appears in \\citet{1997AJ....113.1445W} in the context of small bodies near a binary stellar system, but the coordinates used there were the semi-major axis and the inclination of the orbit, which are less relevant for the kinematic signature explored here. The stability maps elucidate how the velocity phase space populated by stable orbits evolves with integration time, and how it varies with distance from the binary. The maps shows how the velocity distribution function $f(v)$ evolves from a smooth isotropic function at a distance $r\\gg a_{\\rm h}$, to a highly anisotropic function, which produces the BBH kinematic signature.  We derive the BBH kinematic signature by integrating orbits of test particles around a massive BBH, i.e. by solving the restricted 3-body problem, for a binary in a circular orbit, embedded in a bulge potential. We take snapshots of the system once it approaches a steady state, and produce maps of the projected distribution of the stars and their velocity distributions. Similar scattering experiments were already done in the past (e.g. \\citealt{hills83}; \\citealt{mikola92}; \\citealt{1996NewA....1...35Q}; \\citealt{2007ApJ...660..546S}). The purpose of these studies was to derive the effect of the stellar background on the BBH merger rate, while our purpose is the reverse, to derive the effects of the BBH on the background stellar population. The latter calculations were already carried out by \\citet{mm01} based on $N$-body simulations. These simulations were limited to $N\\lesssim10^5$ and cover a region more than $10^3$ larger than $a_{\\rm h}$, thus they do not allow accurate mapping of the projected line of sight velocity distributions (LOSVDs) on scales of a few times $a_{\\rm h}$, as done here. Since the background stellar mass within $a_{\\rm h}$ is negligible (see below), and 2-body scattering cross sections are still small, $N$-body simulations can be replaced by the much faster scattering experiments. These experiments allow us to probe the stellar kinematics with a roughly $10^5$ higher resolution, in the relevant region, compared to \\citet{mm01}. In \\S\\ref{sec:methods} we present the method of the calculation, and in \\S\\ref{sec:results} we present results on the phase space stability regions, maps of the projected velocity distribution moments, velocity moments along some slit positions, various LOSVDs, and also stellar density profiles and projected density maps. The results are discussed in \\S\\ref{sec:discussion}, and summarized in \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} Orbits in the restricted three body problem are notoriously complex (though some insight can be gained from stability maps). Here we exploit this property to derive the signature of a BBH on the nearby stellar kinematics, once unstable orbits are gone. The fraction of velocity phase space populated by stable orbits decreases inwards. The stars ejected by unstable orbits will leave behind a light deficiency, which was suggested to explain the core structure of massive ellipticals. The remaining stars are on significantly anisotropic orbits, characterized by the following properties: \\begin{enumerate} \\item Tangential orbits dominate, mostly retrograde at $r<5$. \\item Increased $\\sigma$, due to the elimination of low tangential velocity orbits. \\item A drop in $\\sigma$ at $r<5$, as most orbits become retrograde. \\end{enumerate} These properties lead to a specific signature on the LOSVD moments on scales as large as 5--10$a_{\\rm h}$, which may be resolved in nearby galaxies. The detection of these kinematic features may indicate the presence of a BBH currently, or a relaxation time ago, beyond which the kinematic signature is erased. If the core structure is a signature of a BBH phase in the past, some BBH kinematic signature may remain on the core scale as well."
    },
    "1005/1005.1639_arXiv.txt": {
        "abstract": "We investigate the characteristics of young ($<$~20~Myr) and bright (L$_X$ $>$ 1 x 10$^{36}$ erg s$^{-1}$) high-mass X-ray binaries (HMXBs) and find the population to be strongly metallicity dependent. We separate the model populations among two distinct formation pathways: (1) systems undergoing active Roche lobe overflow (RLO), and (2) wind accretion systems with donors in the (super)giant (SG) stage, which we find to dominate the HMXB population. We find metallicity to primarily affect the number of systems which move through each formation pathway, rather than the observable parameters of systems which move through each individual pathway. We discuss the most important model parameters affecting the HMXB population at both low and high metallicities. Using these results, we show that (1) the population of ultra-luminous X-ray sources can be consistently described by very bright HMXBs which undergo stable RLO with mild super-Eddington accretion and (2) the HMXB population of the bright starburst galaxy NGC~1569 is likely dominated by one extremely metal-poor starburst cluster. ",
        "introduction": "High-mass X-Ray binaries (HMXBs) form an important observational tool in studies of star-forming regions, due to both their relatively high X-Ray luminosity and age-specificity, which make them good indicators of recent star formation \\citep{2003MNRAS.339..793G}. In extragalactic environments, often only \\emph{bright} HMXBs (here defined as L$_X$ $>$ 1 x 10$^{36}$ erg s$^{-1}$) are detectable. This may significantly affect not only the absolute number of sources, but also the time evolution and observable characteristics of discoverable HMXBs.  Recent studies have noted an inverse correlation between starburst metallicity and the number of observable HMXBs produced by each starburst. \\citet{2004ApJ...609..133M} find the Small Magellanic Cloud (SMC) to be overabundant in HMXBs by a factor of $\\sim$50 compared to the Milky Way - though they note that several factors, including (1) incompleteness in Milky Way observation, (2) recent star formation in the SMC, and (3) the large number of Be-HMXBs observed in the SMC, greatly contribute to this result. Using population synthesis techniques to model continuous star formation activity, \\citet{2006MNRAS.370.2079D} found a factor of $\\sim$3 increase in the total number of HMXBs with L$_X$~$>$~1~x~10$^{32}$~erg~s$^{-1}$ when the metallicity is decreased from solar (Z=Z$_\\odot$) top SMC (Z=0.2~Z$_\\odot$).  Within this population of bright extragalactic HMXBs exists a subpopulation of ultra-luminous X-ray sources (ULX; defined here as L$_{X}$ $>$ 1~x~10$^{39}$~erg~s$^{-1}$, the Eddington limit for a $\\simeq$ 7 M$_\\odot$ black hole), the exact nature of which is not well understood \\citep{1997ApJ...478..542F, 2000ApJ...535..632M, 2001ApJ...554.1035F, 2004PThPS.155...27M, 2004ApJS..154..519S}. Because some of these systems have isotropic luminosities above the Eddington luminosity for black holes (BHs) thought to be produced via single star evolution, much work has centered on the possibility that ULXs contain an intermediate-mass BH (IMBH) \\citep{1999ApJ...519...89C, 2004cbhg.symp...37V, 2008ApJ...688.1235M}. Other arguments point to ULXs as a natural extension of the HMXB population \\citep{2003MNRAS.339..793G}. These models may bypass the strict isotropic Eddington limit by invoking photon bubbles which can escape from a ``leaky disk\" \\citep{2002ApJ...568L..97B}, a non-spherical component of the X-Ray luminosity (beaming) \\citep{2001ApJ...552L.109K} or an optically thin accretion disk \\citep{1988ApJ...332..646A, 2003ApJ...597..780E}. These stellar mass ULXs may achieve super-Eddington luminosities through either strong wind accretion from supergiant donors \\citep{2005MNRAS.356...12S} or thermal timescale Roche-lobe overflow (RLO) \\citep{2001ApJ...552L.109K}.  Observations have found a correlation between ULXs and young starburst activity \\citep{2004ApJ...601L.143I}. The most recent observations have suggested that the number of ULX sources peak between 10-20 Myr after star formation has ended \\citep{2009ApJ...703..159S}. These observations have also found ULX to be most prevelant in low metallicity environments \\citep{2002astro.ph..2488P, 2004ApJ...603..523Z, 2008ApJ...684..282S}. Several models have been posited to explain the metallicity dependence of ULX sources. \\citet{2009MNRAS.395L..71M}, and later \\citet{2009MNRAS.400..677Z} have employed the results of evolutionary calculations that indicate an enhanced formation of high mass black holes ($\\sim$ 30 M$_\\odot$) at low metallicities \\citep{2006A&A...452..295E,2010arXiv1004.0386B}. The large masses of these BHs provide them with Eddington limits which reach typical ULX luminosities ($\\sim$ 10$^{40}$~erg~s$^{-1}$). These massive stellar  black holes are expected to form through direct stellar collapse witn no associated natal kick (e.g., \\citet{2001ApJ...554..548F, 2003Sci...300.1119M}). However, \\citet{2002ApJ...577..710Z} found a moderate displacement between ULX sources and young stellar clusters in the Antennae galaxies, suggesting that the ULX population received a natal kick, indicating that these systems either did not undergo direct collapse, or that direct collapse events are still associated with a small to moderate natal kick.  Seeking to bridge the gap between theoretical predictions and observation, in this \\emph{paper} we analyze the formation of bright HMXBs in the first 20 Myr after starburst activity at several different metallicities. We show that, while our models can produce bright HMXBs at all metallicities, the number of systems, time evolution, and orbital period distribution are strongly metallicity dependent. In particular, while our research supports the observation that ULX are preferentially formed in low-metallicity environments, we find that this relation is not related to the formation of large ($\\sim$~30~M$_\\odot$) direct-collapse BHs at low metallicities, but instead to the dynamics of the pathway which allows RLO onto black holes of $\\sim$10-15 M$_\\odot$ BHs, with mild super-Eddington luminosities.  In Section~\\ref{modeling} we describe our simulation codes and modeling assumptions. In Sections~\\ref{sec:results} and \\ref{sec:pathways} we analyze our results on the properties of bright HMXBs and discuss their physical origin in the context of their evolutionary formation pathways. In Section~\\ref{theorycompare} we discuss the relations between our results and previous work, specifically noting the results for the ULX population in Section~\\ref{subsec:ulx}. Further applying our results to observed systems, in Section~\\ref{subsec:ngc1569} we discuss the production of HMXBs in the starburst galaxy NGC 1569 in the context of the ages and metallicities of the known super star clusters. Finally, we summarize our conclusions in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} We have studied the young, extragalactic, HMXB population and determined that it is composed of systems moving through two distinct binary pathways: mass transfer systems main sequence donors and wind accretion systems containing (super)giant donors. Metallicity greatly affects the fraction of systems which move through each pathway, but generally does not change the characteristics of binaries moving through a specific pathway. From these results, we have made three observationally testable claims: (1) the majority of the ULX population can be explained as an extension of the stellar HMXB population, (2) ULXs should be preferentially found in low metallicity clusters - while the ULXs found in higher metallicity regions should be younger and have larger orbital periods, and (3) the starburst galaxy NGC 1569 is dominated by contributions from Cluster A, which likely provides approximately 90\\% of the observed HMXB population. Detailed modeling of both theoretical phenomena, such as the effect of non-spherical winds falls outside the scope of this paper, but is currently ongoing. New observations of starbursts in low metallicity environments are also underway, and will have the power to quantitatively test our predictions and determine the important role that metallicity plays in HMXB evolution."
    },
    "1005/1005.1125_arXiv.txt": {
        "abstract": "This paper presents the first of two analyses about the influence of environment on the formation and evolution of galaxies observed in the nearby universe. For our study, we used three different samples representing different density environments: galaxies in Compact Groups (HCGs), Isolated Pairs of Galaxies (KPGs),  and Isolated Galaxies (KIGs), which were taken as references. Using both characteristic isophotal parameters and evidence of asymmetries in the optical and the near-infrared, we are able to establish differences in the characteristics of galaxies with different morphologies in different environments, allowing us to better understand their different formation histories. In this first paper we present the isophotal and asymmetry analyses of a sample of 214 galaxies in different environments observed in the optical ($V$ and $I$ images). For each galaxy, we have determined different characteristic isophotal parameters and $V - I$ color profiles, as a function of semi-major axis, and performed a full asymmetry analysis in residual images using the $V$ filter. Evidence of asymmetry in the optical is almost missing in the KIG sample, and significantly more common in the KPG than in the HCG samples. Our isophotal analysis suggests that the stellar populations in the HCG galaxies are older and more dynamically relaxed than in the KPG. The HCG galaxies seem to be at a more advanced stage of interaction than the KPGs. One possible explanation is that these structures formed at different epochs: compact groups of galaxies would have formed before close pairs of galaxies, which only began interacting recently. However, similarities in the formation process of galaxies with same morphology suggest CGs and close pairs of galaxies share similar conditions; they are new structures forming relatively late in low-density environments. ",
        "introduction": "Although it is recognized that the environment of galaxies plays an important role in their formation and evolution, the mechanisms responsible for such processes, the details on how they apply, and the time-scales on which they are effective are still largely unknown. For example, in Compact Groups of Galaxies (CGs) we have recently shown \\citep{cp07} that mergers and tidal interactions are two important mechanisms driving the morphological evolution of galaxies in these systems. We have also found that many of the ongoing merger events were possibly happening without gas, a phenomenon known in the literature as a ``dry merger'' \\citep{vandokkum05}. According to the dry merger hypothesis, elliptical galaxies are generally formed by the merger of bulge-dominated galaxies, not from the merger of spiral-like galaxies. This is fully consistent with the CG environment where early-type galaxies constitute the dominant population \\citep{coziol04}. Finding evidence of dry mergers in CGs is important, because it suggests that these systems are obviously not in a dynamically stable state. It also suggests that since the merging process of the galaxies is not yet complete, these systems cannot be as old as that believed based on the absence of standard observational merging evidence, like luminous active galactic nuclei (AGNs), recent star formation events, or post-burst stellar populations in evolved galaxies.  What is missing in CGs is a time scale for the evolution process of the galaxies. Was the evolution of galaxies accelerated in the group environment? Is the dry merger the result of such evolution? Is the dry merger limited only to dense environments? What is the role of the potential of the group in the disappearance of the gas? How fast was the gas exhausted or consumed? Did it burn rapidly forming the bulges of numerous early-type galaxies, was it lost feeding a black hole, or was it mostly ripped off of the galaxies and lost to the intergalactic medium?  In order to find some answers to the above questions, we have decided important to extend our study of characteristic isophotal parameters and asymmetry to two different structures having lower spatial density than CGs: isolated galaxies and isolated pairs of galaxies. Isolated galaxies are considered to have a low probability of interaction with another galaxy of similar mass over a Hubble time \\citep{vettolani86}. Consequently, a sample of isolated galaxies can be treated as a set of ``comparison objects'', free during most of their lifetime from environmental effects. Isolated galaxies are uncommon in the universe where most of the galaxies tend to be clustered in groups, as shown by Tully (1987). The reason for their existence, therefore, may be an interesting subject of study by itself. Isolated pairs of galaxies ar in the next level of galaxy density. In the nearby universe, these systems are also rare and consequently their history is not well documented. Many questions still need to be answered. How long have these galaxies been interacting? Are they engaged in first encounters or did they interact multiple times before with their companions? Are these transient phenomena (high-energy orbits) or merging encounters ( low-energy or decaying orbits)?  For our study, we used three well-defined environment samples: the Catalog of Isolated Galaxies (KIGs), from Karachentseva (1973), the list of Isolated Pairs of Galaxies (KPGs), as compiled by Karachentsev (1972), and the Catalog of Compact Groups of Galaxies (HCGs) from Hickson (1982). Our analysis is based on the application of two independent methods: the fitting of elliptical ellipses on the isophotal levels of the galaxies and the determination of their asymmetry \\citep{cp07}. We present the characteristics of the observed samples in Section~2 and describe our observations and the reduction process in Section~3. In Section~4, we explain the methods used for our analysis. The surface photometry profiles, the color maps and asymmetrical images of the observed galaxies, and the results of nonparametric statistical tests used to establish the level of significance of the differences observed are presented in Section~5. Finally, we discuss our results and give our conclusions in Section~6. Our analysis for the near-infrared will follow in an accompanying paper \\citep{plauchu10}. ",
        "conclusions": "Through our analysis we have found that galaxies in close pairs show more frequent signs of interactions at a higher level than those in compact groups. This may seem somewhat contradictory. If interaction between galaxies is favored in high-density environments with low velocity dispersion, should we not expect evidence for such processes to be more obvious in multiple systems like CGs? A possible answer to this apparent contradiction can be found in our isophotal analysis. Indeed, we have seen that the HCG galaxies tend to be redder in their central part and less compact in their periphery than the KPG galaxies, which is consistent with older central stellar populations and dynamically more relaxed orbits as a whole in the HCG galaxies than in the KPG galaxies. These observations, together with the presence of asymmetries at a lower level  in the HCG galaxies, suggest CGs are found in a more advanced stage of interaction than pairs of galaxies. One possible explanation is that these structures formed at different epochs: CGs would have formed in the recent past, while close pairs would have formed even more recently.  The alternative interpretation is to assume that the evolution of galaxies is accelerated in CGs: the galaxies in CGs formed at the same time as those in close pairs, but they evolved faster due to multiple interactions. However, based on our observations, such an alternative seems less probable. In particular, we observe similar properties for the E$-$S0 in the HCGs and KPGs which suggest similar formation mechanisms. The higher number of such galaxies in the HCGs (see Figure~\\ref{comp}), therefore, can only be the result of originally higher matter density: in denser regions, a high number of galaxies are formed, which can eventually interact to build larger and more complex structures like CGs, while in less dense environments, a few galaxies are formed and it can take longer for these galaxies to interact with neighbors.  On the other hand, the fact that many S0 galaxies can also be found among the KPGs suggests their environment must have some level of similarity with that of the HCGs. The common property is that both systems are examples of structures forming in relatively low density environments; that is, both form relatively late at the periphery of large-scale structures.  The cosmological model that better fits our observations is one where the formation of structures is a biased process. As a consequence, it is expected that massive structures, which formed in originally denser regions, must assemble their components at earlier epochs than less massive ones. If we also assume the formation process of structures to be continuous in time, then we must now expect to observe smaller mass structures like CGs and pairs of galaxies to form at the periphery of the larger-scale structures."
    },
    "1005/1005.3911_arXiv.txt": {
        "abstract": "Anomalous X-ray Pulsars and Soft Gamma-Ray Repeaters have been generally recognized as neutron stars with super strong magnetic fields, namely ``magnetars\". The ``magnetars\" manifest that the luminosity in X-ray band are larger than the rotational energy loss rate, i.e. $L_{X}>\\dot {E}_{\\rm rot}$, and then the radiation energy is coming from the energy of magnetic field. Here it is argued that magnetars may not really exist. Some X-ray and radio observational results are contradicted with the magnetar model.  (1) The X-ray luminosity of PSR J1852+0040 is much larger than the rotational energy loss rate ($L_{X}/\\dot {E}_{\\rm rot}\\simeq 18 )$, but the magnetic field is just $3.1\\times 10^{11}$\\,G. Does this X-ray radiation energy come from the magnetic field?  (2) In contrast to the above, the magnetic fields of radio pulsars J1847-0130 and PSR J1718-3718 are higher than that of AXP 1E 2259+586, why is the radiation energy of those two radio pulsars still coming from rotational energy? Furthermore, the magnetic field of the newly discovered SGR 0418+5729 with the lowest magnetic field is 3.0\\,e13\\,G, lower than the critical magnetic field $B_{\\rm C}=4.414$\\,e13\\,G) (Esposito et al. 2010).  (3) Some ``magnetars\" also emit normal transient radio pulses, what is the essential difference between radio pulsars and the ``magnetars\"?  The observational fact arguments will be presented at first, then we discuss in what situation the conventional method to obtain magnetic field could not be correct. ",
        "introduction": "Anomalous X-ray Pulsars (AXPs) and Soft Gamma-Ray Repeaters (SGRs) have similar characteristics (see, e.g. Mereghetti et al., 2009; Kaspi, 2007; 2009). Timing analysis gives no evidence of orbit motion, no evidence to show AXPs and SGRs in binary systems. Hulleman et al. (2000) found that the optical emission from the AXP (4U 0124+61) is too faint to admit a large accretion disk.  AXPs and SGRs are recognized as the new classes of young objects ($P/(2 \\dot P) \\sim 10^3 - 10^5 $\\,yr), along with other young neutron star families, and the magnetic fields are super strong ($5.9\\times10^{13}$\\,G to $10^{15}$\\,G). They do not manifest themselves as radio pulsars, like the dim isolated neutron stars (DINs) and the center compact objects (CCOs), which suggests alternative evolutionary paths for young neutron stars (Ertan et al. 2009).  The most peculiar observed fact is that all AXPs and SGRs have stable persistent pulsed X-ray emission with the luminosity of $L_{X} \\sim 10^{34} - 10^{36}$\\,ergs\\,s$^{-1}$, well in excess of the spin down energy of these sources, i.e., $L_{X}>\\dot{E}_{\\rm rot}$.  A very important question is then: what does the radiation energy come from?  To answer this question, some models are suggested. Meanwhile, the ``magnetar\" model is widely accepted, which is isolated neutron stars with super strong magnetic fields (Duncan \\& Thompson 1992; Thompson \\& Doncan 1995), and the radiation energy is coming from the magnetic field decay (Thompson \\& Doncan 1996).  For AXPs and SGRs, the widely accepted assumption is that the rotational energy is carried out by the dipole magnetic radiation. This is just the fact what we have observed for radio pulsars. So radio pulsars are named as ``spin-down powered\" neutron stars. Unfortunately for AXPs and SGRs, no direct link between rotational energy and the dipole magnetic radiation has been found so far, thus any calculations for the value of magnetic field obtained from $P$ and $\\dot{P}$ could not be reliable.  Furthermore, there are some observations that challenge the ``magnetar\" model.  (1) X-ray observations show at least one CCO, PSR J1852+0400, in the supernova remnants Kes 79, shows ``anti-magnetar\" properties: the field is low ($3.1\\times 10^{10}$\\,G), but X-ray emission is larger than the spin down energy, $L_{X}>\\dot{E}_{\\rm rot}$. What does the radiation energy come from? Does it also come from the magnetic field?  (2) Some authors argue that owing to the super strong magnetic field, so AXPs and SGRs are different from the normal radio pulsars: magnetars can not radiate in the radio band; and radio pulsars can not have super strong magnetic field. But recent observations show that some AXPs: XTE J1810-197 (Camilo et al. 2007) and 1E 1547.0-5408 (Camilo et al. 2008) can have normal radio pulse emission. On the other hand, some normal radio pulsars PSR J1847-0130 and PSR J1718-3718 do have super strong magnetic field , which are larger than that of one of the AXPs. In this case, can we say that the differences between ``magnetars\" and radio pulsars are coming from the difference of the magnetic field? If not, what is the real difference?  In $\\S$ 2, we present the anti-magnetar observations in both X-ray and radio bands.  In $\\S$3, it is argued that the magnetic field values of ``magnetsrs\" are incorrect.  The conclusion and discussion are presented in $\\S 4$. ",
        "conclusions": "From the discussions above, our conclusions are drawn as follows:  1. An anti-magnetar (PSR J1852+0040) shows that the X-ray luminosity is larger than the rotational energy loss rate, that is $L_{X}/\\dot{E}>1$. The magnetic field is $3.1\\times 10^{10}$\\,G. Where does the radiation energy come from? Can we say that it is also coming from the magnetic energy?  We argue that in the case of $L_{X}/\\dot{E}>1$, using $P$ and $\\dot{P}$ to get the magnetic field is incorrect, flow-out particles and the parallel magnetic component, $\\mu_{\\parallel}$, should be taken into account, which contributes to the rotational energy loss (Xu \\& Qiao 2001).  2. The differences of ``magnetars\" and the radio pulsars are not from the difference of the magnetic field strength. The transient pulsed radio emission from ``magnetars\" is a kind of typical radio emission as similar as the one observed for normal radio pulsars. The flat spectrum of magnetars is not different from normal radio pulsars. An important thing is that, to show the distinct difference between the transient pulsed radio emission from the magnetars (that is related to the outburst) and the normal radio pulsars.  3. Many authors show that how to produce the observed phenomena of magnetars, but never to show the differences between normal radio pulsar and magnetars. For example, why magnetars have high energy radiations but nor for the radio pulsar PSR J1847-0130.  We argue here that a acceptable theory, related to observations of manetars, should show the difference between magnetars and other objects.  4. Producing the very strong outbursts is a key question for any model about the AXPs \\& SGRs. If the very high energy of the outburst does not come from the magnetic field, where oes the energy come from? Only accretion may not satisfy this request. Some possible quark star models should be taken into account, e.g. Cheng \\& Dai (1998), Cheng et al. (2000); Xu, Tao \\& Yang (2006)."
    },
    "1005/1005.3066_arXiv.txt": {
        "abstract": "The 2009 November outburst of the neutron star X-ray binary Aquila X-1 was observed with unprecedented radio coverage and simultaneous pointed X-ray observations, tracing the radio emission around the full X-ray hysteresis loop of the outburst for the first time.  We use these data to discuss the disc-jet coupling, finding the radio emission to be consistent with being triggered at state transitions, both from the hard to the soft spectral state and vice versa.  Our data appear to confirm previous suggestions of radio quenching in the soft state above a threshold X-ray luminosity of $\\sim10$\\% of the Eddington luminosity.  We also present the first detections of Aql X-1 with Very Long Baseline Interferometry (VLBI), showing that any extended emission is relatively diffuse, and consistent with steady jets rather than arising from discrete, compact knots.  In all cases where multi-frequency data were available, the source radio spectrum is consistent with being flat or slightly inverted, suggesting that the internal shock mechanism that is believed to produce optically thin transient radio ejecta in black hole X-ray binaries is not active in Aql X-1. ",
        "introduction": "Multi-wavelength observations of accreting stellar-mass compact objects have revealed a fundamental coupling between the processes of accretion and ejection.  As an X-ray binary (XRB) evolves through a set of characteristic X-ray spectral and variability states over the course of an outburst, the radio emission, assumed to arise from jets launched from the inner regions of the system, is also seen to evolve. \\citet{Fen04} proposed a phenomenological model for this accretion-ejection (disc-jet) coupling in the relatively well-studied black hole (BH) XRBs.  In this model, compact steady jets are present in the hard X-ray spectral state, with the jet power scaling with X-ray luminosity.  The jet velocity increases when the source makes a transition to a softer spectral state, leading to internal shocks within the flow \\citep[e.g.][]{Kai00}, observed as discrete knots of radio emission that move outwards at relativistic speeds.  The compact jets then switch off.  The system eventually makes a transition back to a hard spectral state, whereupon the core jets are re-established, and the source fades back into quiescence.  This second state transition tends to occur at a lower luminosity than the original transition from the hard to the soft state, implying a hysteresis in the outburst cycle \\citep{Mac03}.  Compact jets are also inferred to exist in neutron star (NS) systems, from brightness temperature arguments, the observed flat-spectrum radio emission, and the detection of a jet break in broadband spectra \\citep{Mig06b,Mig10}.  Rapid energy transfer from the core to detached radio lobes \\citep{Fom01,Fen04} also argues strongly for the existence of jets in NS systems, although their compact jets have not been directly imaged, as in BH systems \\citep{Dha00,Sti01}.  Since NS XRBs are typically fainter radio emitters than their BH counterparts at the same X-ray luminosity \\citep{Fen01,Mig06}, the nature of the disc-jet coupling in NS systems is consequently less well understood. \\citet{Mig06} made a systematic study of radio emission from NS XRBs, finding evidence for steady jets in hard state systems at low luminosities ($<1$\\% of the Eddington luminosity, $L_{\\rm Edd}$) and transient jets in outbursting sources close to $L_{\\rm Edd}$, as seen in BH systems but without complete suppression of the radio emission in soft states.  The correlation between X-ray and radio emission in hard states was steeper in NS systems \\citep{Mig06} than the analogous correlation in BH systems \\citep{Gal03}, implying a different coupling between X-ray luminosity and mass accretion rate, although the jet power for a given accretion rate is similar in both classes of system \\citep{Koe06}.  Understanding the similarities and differences between the disc-jet coupling in BH and NS systems is crucial in determining the role played by the depth of the gravitational potential well, the stellar surface, and any stellar magnetic field in the process of jet formation.  \\subsection{Aql X-1}  Aql X-1 is a recurrent transient XRB which undergoes outbursts every $\\sim300$\\,d.  The accretor is a confirmed NS \\citep{Koy81}, in a $\\sim19$-h orbit \\citep{Che91,Sha98,Wel00} with a K7V companion \\citep{Che99}.  Its X-ray spectral and timing behavior classify it as an atoll source \\citep{Rei00}.  The distance, determined from the luminosity of its Type I X-ray bursts, is in the range 4.4--5.9\\,kpc \\citep{Jon04}.  Despite its frequent outbursts, there are few reported radio detections of Aql X-1, likely owing to the faintness of atoll sources in the radio band \\citep{Mig06}.  \\citet{Tud09} analyzed all publicly-available archival data from the Very Large Array (VLA), taken exclusively during X-ray outbursts, between 1986 and 2005. They detected the source at 11 epochs, with a maximum radio flux density of 0.4\\,mJy.  In all cases where multi-frequency observations were available, the radio spectrum was inverted ($\\alpha>0$, where flux density $S_{\\nu}$ scales with frequency $\\nu$ as $S_{\\nu}\\propto\\nu^{\\alpha}$).  They also found tentative evidence for quenching of the radio emission above a certain X-ray luminosity.  Here, we extend on the work of \\citet{Tud09}, presenting the most complete radio coverage of an outburst of Aql X-1 obtained to date, to elucidate the nature of the radio/X-ray coupling in this source. ",
        "conclusions": "We have obtained unprecedented radio coverage of an outburst of the NS XRB Aql X-1, demonstrating that the radio emission is consistent with being activated by both transitions from a hard state to a soft state and by the reverse transition at lower X-ray luminosity.  Our data appear to confirm previous suggestions of quenched radio emission above a certain X-ray luminosity, corresponding to $\\sim10$\\% of the Eddington luminosity. We have for the first time detected the source with VLBI, which together with the measured brightness temperature, radio spectra and integrated flux density observed by the VLA demonstrates that the radio emission in this source is consistent with the compact, partially self-absorbed steady jets seen in the hard states of BH XRBs.  NSs thus appear equally capable of launching jets as BHs, but the absence of bright, optically thin, relativistically moving knots may suggest a fundamental difference in the jet formation process between NS and BH systems, possibly hinting at an additional role from the ergosphere."
    },
    "1005/1005.2195_arXiv.txt": {
        "abstract": "We present the results of a program of $K-$ and $K_S-$band imaging of a sample of 2Jy radio galaxies with redshifts $0.03 \\lta z \\lta 0.5$, for which the host galaxy morphologies and structural parameters (effective radius, S\\'ersic index and unresolved nuclear point source contribution) have been determined using \\textsc{galfit}.  Two-thirds of our sample are best modelled as being hosted by massive elliptical galaxies with S\\'ersic indices of $n=4-6$, with the remainder being better suited either by a mixture of morphological components (usually a bulge plus a small, less luminous, disk component) or by more disky galaxy models with $n=1-2$. Our measured galaxy sizes are generally in very good agreement with other imaging programs, both space- and ground-based.  We also determine a slightly higher average nuclear point source contribution than similar HST-based programs. This is due to our inability to separate the AGN emission from compact circum-nuclear stellar emission, but does not bias our modelling of the remainder of the host galaxies and our results remain robust.  We also observe that roughly half of the objects in our sample are either undergoing major or minor merger activity or are clearly morphologically disturbed. ",
        "introduction": "Over the last decade, it has become increasingly apparent that not only do supermassive black holes reside in the centres of all massive galactic bulges, but also that their ongoing evolution is strongly coupled, despite nine orders of magnitude difference between the relevant spatial scales of the black hole and its galactic host (Kormendy 1993; Magorrian et al 1998; Ferrarese \\& Merritt 2000; Gebhardt et al 2000; Merritt \\& Ferrarese 2001; Cattaneo \\& Bernardi 2003; Marconi \\& Hunt 2003; H\\\"aring \\& Rix 2004). Regardless of the physical driving mechanisms for these correlations, it is certain that understanding the periods of AGN activity, when the central black hole is most able to grow via accretion, are of crucial importance for building up a cogent theoretical framework for galaxy evolution in general (e.g. Yu \\& Tremaine 2002, Shankar et al 2004, Marconi et al 2004 and references therein). This problem poses two central, interlinked questions: what is the nature of an active galaxy, and how do they differ from the inactive population of galaxies from which they are (presumably) drawn?  Extragalactic radio sources play a particularly important role in such studies. Historically, they have been amongst the best studied distant AGN, due to the ease with which large samples of such objects could be assembled out to large redshifts. Under orientation-based unification schemes for radio galaxies and radio loud quasars (e.g. Barthel 1989; Urry \\& Padovani 1995), they provided a perfect laboratory for studying the host galaxies of radio loud AGN without the complicating presence of an unobscured bright quasar nucleus.  Additionally, the presence of extensive powerful radio emission directly links the behaviour of these systems from the small scales of the central accreting black hole out to many hundreds or thousands of kiloparsecs into the surrounding intergalactic medium. Early studies of the proto-typical 3CR sources suggested that the host galaxies of these radio sources were members of a remarkably uniform population, and seemed to be ``standard candles'' to within $\\sim 0.3$mag (Lilly \\& Longair 1984), with characteristic effective radii of up to 15kpc and IR magnitudes dominated by emission from an old stellar population formed at $z > 5$ (Best et al 1998). However, subsequent revisions to the most widely accepted cosmological model (e.g. Hinshaw et al 2009) and an improved understanding of structure formation in the universe (e.g. Springel et al 2005) add considerable complexity to this picture.  While a radio source {\\it may} undergo several active phases (e.g. Haehnelt \\& Rees  1993; Saikia \\& Jamrozy 2009 and references therein), this does not imply that any given radio source would necessarily have been active in the past, or will be so again in the future.  This opens up the scenario whereby radio galaxies at different epochs could be drawn from quite different galaxy populations: might this indeed be the case?  The tightness of the  $K-z$ relation (e.g. Inskip et al 2002; Willott et al 2003; Rocca-Volmerange et al 2004) and the total  near-IR luminosities of their host galaxies (Seymour et al 2007) suggest that radio sources are hosted by a relatively homogeneous, high mass  (up to $\\sim 10^{12} \\rm M_{\\odot}$) population of galaxies over a wide redshift range. Even at higher redshifts, while morphological studies have shown that some sources are still undergoing active assembly, radio source host galaxies at redshifts of $z \\sim 2-3$ are in general similar in size to those nearby (e.g. van Breugel et al 1998, Pentericci et al 2001). However, the increased luminosity of high-redshift radio galaxies relative to passive evolution models goes against the simplest picture of closed-box evolution of a single population of galaxies.  The high mass of radio source host galaxies at all epochs provides an additional constraint.  Since $z \\sim 1$, the most massive members of the elliptical galaxy population have shown little evolution in mass (e.g. Heavens et al 2004; Thomas et al 2005). Although subsequent episodes of AGN activity in massive galaxies could be triggered by minor mergers or galaxy interactions which lead to little additional change in galaxy mass, the links between periods of black hole accretion and galaxy growth (as evidenced by the black hole mass vs. bulge mass relation) limits the ability of such objects to repeatedly host powerful, high accretion-rate AGN in more recent cosmic epochs. On top of this, it has been theorised that the very presence of a powerful radio source and/or AGN can limit the potential for further galaxy growth, via heating of the local environment (e.g. Fabian et al 2000; McNamara \\& Nulsen 2007; Croton \\& Farrar 2008; De Young 2010).  Equally, the actively accreting lower-redshift systems have likely been assembled from their less-massive constituent progenitor objects too recently to have hosted powerful AGN at high redshifts.  Similarly, at different redshifts the powerful radio source population samples different environments (e.g. Lilly \\& Prestage 1987; Yates et al 1989; Hill \\& Lilly 1991; Best 2000), with a tendency towards more highly clustered objects at higher redshifts. It is now clear that galaxies are subject to so-called ``cosmic downsizing'', by which both star formation (e.g. Cowie et al 1996; Juneau et al 2005) and AGN activity (e.g. Granato et al 2004; De Lucia et al 2006, Croton et al 2006; Merloni \\& Heinz 2008; Martini et al 2009; Shen 2009) become increasingly more prevalent in lower-mass  objects at later cosmic epochs, and in less dense environments.  The inherent changes in the host galaxy population highlight the essential need to effectively characterise their properties at all redshifts. Without a clear handle on what a radio galaxy {\\it is} at any given epoch, it is practically impossible to determine how it may or may not differ from the underlying inactive galaxy population, or from the population of radio-quiet AGN.  What are the characteristic sizes of powerful radio galaxies, and how does this parameter vary with redshift and/or radio power? Is the scatter in the $K-z$ relation solely due to variations in galaxy mass, or can it be accounted for at least in part by variations in the host galaxy stellar populations?  Do the host galaxies show signs of recent merger activity, and how does the proportion of merging/disturbed sources compare with that of normal galaxies? Is there any dependence of galaxy morphology on local environment, or on the radio source/AGN properties? With such physical characteristics of the host galaxies pinned down, deeper questions may then be posed on the nature of {\\it how} AGN and radio source activity are triggered, e.g.: whether triggering occurs due to major or minor interactions or the more gradual infall of material; the prevalence of different triggering modes; the details of the timeline of events involved and any timelags between star formation and AGN/radio source activity; and the links between such processes and the evolution and  observed morphological characteristics of the host galaxies.  In order to investigate these issues, we turn not to the 3CR sample, but instead to the similarly well studied 2Jy sample of southern radio sources at $z < 0.5$ (Tadhunter et al 1993, Morganti et al 1993). This sample is unique in the sense that it is the {\\it only} major sample of powerful radio galaxies for which complete, high quality optical continuum and emission line spectra, as well as radio imaging data, exist for all the objects, alongside near- and mid-IR imaging and spectroscopy for most sources.  One particular advantage of this sample is that the combination of the mid-IR data and optical spectra allows us to separately determine the extent of both obscured and unobscured star formation activity within the host galaxies, as well as providing two independent measures of the degree of nuclear activity (via the 24$\\mu$m emission and the optical line emission). Our current dataset consists of ground-based near-IR imaging of the majority of the sample objects.  These data will be used to obtain accurate photometric magnitudes, and to thoroughly characterise the host galaxy structural properties as observed in this wavelength regime (galaxy profile, effective radius and the extent of any nuclear point source contamination from the AGN), and how they compare to other radio galaxy samples.  We will also make a preliminary assessment of the prevalance of visible merger activity within the sample (either ongoing or having occured in the recent past); this issue is expanded on in our analysis of optical imaging data for the 2Jy sample (Ramos Almeida et al in prep, hereafter RA10).  Having characterised the morphological properties of the host galaxies, in paper 2 of this series we will extend our investigation to the more global properties of the sample.  By combining galaxy morphologies with our data on the host galaxy stellar populations, the results of the current paper will be used to investigate the nature of the scatter in the $K-z$ relation, and how well it can be accounted for in terms of galaxy mass and recent star formation activity. We will also use literature data for other radio galaxy surveys at higher redshift to examine how redshift and radio power impact on the observed trends. Finally, together with the results of RA10, we will assess in paper 2 how closely morphological signatures of merger activity correlate with the degree of recent star formation, the host galaxy morphologies, and the size/age of the radio source, helping us to build a clearer picture of the variety of different physical modes of AGN triggering in these sources. Further to this work, these data will also be used in conjunction with the existing optical data (RA10) to also constrain the local environments of the host galaxies (Ramos Almeida et al, in prep), and to comprehensively assess the impact of environment on the properties of powerful radio sources and their host galaxies.  In section 2, we provide details of our observed sample. Our observations and data reduction processes are described in section 3. In section 4, we present our observational results and our 2-d modelling of the host galaxies, while our discussion and conclusions are presented in sections 5 and 6 respectively.  A more extensive analysis of this dataset is postponed to a second paper in this series.  Throughout this paper, we assume cosmological parameters of $\\Omega_0 = 0.27$, $\\Omega_{\\Lambda} = 0.73$ and $H_0 = 71 \\rm{km\\,s^{-1}\\,Mpc^{-1}}$. ",
        "conclusions": "We have presented near infrared imaging observations for 41 radio galaxies of the 2Jy sample with $0.03 \\lta z \\lta 0.5$. The sources lie close to the tracks for passively evolving galaxies formed at high redshift on the infrared Hubble diagram.  Through a combination of aperture photometry and 2-d surface profile modelling using \\textsc{galfit}, several useful galaxy parameters have been quantified for the sample, including the percentage contamination by a nuclear point source, S\\'ersic index and effective radius.  Overall, our structural parameter modelling results are in very good agreement with those in the literature for other radio galaxies at higher and lower redshifts, though a small fraction of the unresolved nuclear emission detected in our objects may be due to unresolved cuspy nuclear starlight profiles rather than AGN emission alone.  Roughly two thirds of the sources in our sample are hosted by massive elliptical galaxies, with a further 20\\% being hosted by galaxies best modelled as bulges with an additional disk-component contributing additional flux to the galaxy. Half are clearly either disturbed or interacting with companion objects.   Except at the lowest redshifts (where the minimum intrinsic radio power required for membership in a flux limited sample is lower), disky-type galaxies rarely host the powerful radio sources of samples such as the 2Jy sample, unless other events (such as a major merger) increase the likelihood of a radio source having been triggered.  Although we see no obvious trends between AGN type, recent/ongoing merger activity and a tendency towards galaxies being best modelled by either multi-component or disky models, the physical links between AGN activity and the host galaxy properties may indeed be very indirect, and subject to time delays between mergers/interactions and any subsequently triggered AGN activity. In the second paper of this series, we investigate the case for such links in greater depth.  There, we extend our investigation of the global properties of this sample, tying the newly-measured structural parameters to the other properties of these systems (including the radio source properties, the ongoing visible and obscured star formation, and the interaction status of the host galaxies). The derived galaxy luminosities, morphological parameters and measures of star formation activity are used to fully characterise the scatter on the infrared Hubble diagram. With the aid of literature data, we will also assess the dependence of the host galaxy characteristics on redshift and radio power.  Finally, the combination of optical (from RA10) and near-IR signatures of recent merger activity will be correlated with the overall galaxy morphologies, star formation histories and radio source properties to highlight the links between radio source triggering/AGN activity and the ongoing evolution of the host galaxies, and to develop a more thorough picture of the diversity of radio source triggering mechanisms, and the timelines/chains of events involved."
    },
    "1005/1005.2689_arXiv.txt": {
        "abstract": "{Detached circumstellar dust shells are detected around three carbon variables using \\textit{Herschel}-PACS. Two of them are already known on the basis of their thermal CO emission and two are visible as extensions in IRAS imaging data. By model fits to the new data sets, physical sizes, expansion timescales, dust temperatures, and more are deduced. A comparison with existing molecular CO material shows a high degree of correlation for TT~Cyg and U~Ant but a few distinct differences with other observables are also found.} ",
        "introduction": "For low- and intermediate-mass stars (of initial mass less than 8~$M_{\\sun}$) mass-loss mainly takes place on the thermally pulsing asymptotic giant branch (AGB) in slow (typically $5 - 25$~km~s$^{-1}$) winds with large mass-loss rates (up to 10$^{-4}$~$M_{\\sun}$~yr$^{-1}$). Whereas the gaseous component of the mass-loss can be observed from the ground (mainly on the basis of molecular emission lines in the mm- and sub-mm radio range), the interesting dusty fraction of the wind is preferably observed with infrared space telescopes (see e.g., the review by Olofsson, \\cite{Olofs03} and references therein). IRAS was the first infrared space observatory to enable the identification of a large number of AGB stars with extended dust shells. Young et al. (\\cite{Young93b}) interpret them in terms of an interaction of the expelled material with the ISM.  In the generally accepted picture, mass-loss increases on average during the evolution along the AGB (Habing \\cite{Habin96}). Nevertheless, observational evidence implies a more episodic mass-loss evolution (van der Veen \\&{} Habing \\cite{VeenH88}, Olofsson et al. \\cite{Olofs88}). Based on IRAS photometry, van der~Veen \\&{} Habing, identified a group of 60~$\\mu$m excess objects with very low temperature dust, Olofsson et al. found -- in the course of a larger CO survey -- stars with obviously detached gas shells. These two indicators of perhaps the same phenomenon were subsequently used to study this interesting evolutionary phase.  Waters et al. (\\cite{Water94}) and Izumiura et al. (\\cite{Izumi95}) reconstructed high resolution images using IRAS data at 60 and 100\\,$\\mu$m and were able to detect a few detached dust shells (e.g., \\object{U~Ant}, \\object{U~Hya}, or \\object{Y~CVn}). Soon after this, with the improved sensitivity of the ISO satellite, Izumiura et al. (\\cite{Izumi96}) were able to image \\object{Y~CVn} at both 90 and 160\\,$\\mu$m. Nevertheless, imaging of xtended dust emission was hampered notoriously by the low spatial resolution that space telescopes of diameter 60 or 80\\,cm could deliver at long wavelengths ($\\ge 60$\\,$\\mu$m) to probe the cold dust of the detached envelopes (compare the Akari and Spitzer results presented in Izumiura \\cite{Izumi09} and Geise et al. \\cite{Geise10}, respectively). Here the 3.5\\,m telescope of the \\textit{Herschel} Space Observatory (Pilbratt et al. \\cite{Pilbr10}) using PACS (Poglitsch et al. \\cite{Pogli10}) is expected to play in a different league.  Over the past two decades, the study of detached shells in the CO radio line emission has turned out to be most fruitful. Single dish surveys (Olofsson \\cite{Olofs96}, and references therein) have revealed a number of these interesting objects, whereas follow-up interferometric maps have uncovered their detailed structure (Lindqvist et al. \\cite{Lindq99}, Olofsson et al. \\cite{Olofs00}). The available high spatial resolution has detected remarkably spherical and very thin CO-line emitting shells, which are indicative of quite short phases of intense mass-loss probably associated with thermal pulses. In the case of TT~Cyg, the authors speak of a high mass-loss phase, lasting only a few hundred years, that occurred 7000 years ago. This basic scenario may be modified by more complex wind--wind interactions (Maercker et al. \\cite{Maerc10}). In  general, the detached molecular shells typically detected are geometrically thin and probably smaller in size than the detached dust shells found by IRAS. It may be that they have a different origin -- episodic mass-loss on the one hand and wind--interstellar medium (ISM) or wind--wind interaction on the other. Objects such as  U Ant (Izumiura et al. 1997) can exhibit both effects. ",
        "conclusions": "All \\textit{Herschel}-PACS maps detect circular rings of emission around the three central stars. These detached shells suggest that the stars underwent an increase in mass-loss for a short period of time. In all cases, we have found that varying dust emission continues to the inner parts of the detached shells. This can be seen from the  intensity profiles in Fig.~\\ref{profiles} and the models. The corresponding timescales are of the order of thousands of years, whereas the age of the outer detached material is some 1000 or even 10\\,000 years. The correlation with the molecular observations is striking in some cases even down to fine details, e.g., the lower shell density of TT~Cyg to the north. More detailed modelling combining all information about the objects, additional observations, and comparison with future \\textit{Herschel} observations of similar objects will definitely provide a clearer understanding of the time evolution of the mass-loss process."
    },
    "1005/1005.5126_arXiv.txt": {
        "abstract": "We measure the rest-frame colors (dust-corrected), infrared luminosities, star formation rates, and stellar masses of 92 galaxies in a Spitzer-selected cluster at $z=1.62$.  By fitting spectral energy distributions (SEDs) to 10-band photometry ($0.4\\mu$m$<\\lambda_{obs}<8\\mu$m) and measuring \\mipsmu\\ fluxes for the 12 spectroscopically confirmed and 80 photometrically selected members, we discover an exceptionally high level of star formation in the cluster core of $\\sim1700$\\msunyr~Mpc$^{-2}$.  The cluster galaxies define a strong blue sequence in $(U-V)$ color and span a range in color.  We identify 17 members with L$_{IR}>10^{11}$~L$_\\odot$, and these IR luminous members follow the same trend of increasing star formation with stellar mass that is observed in the field at $z\\sim2$.  Using rates derived from both the \\mipsmu\\ imaging and SED fitting, we find that the relative fraction of star-forming members triples from the lowest to highest galaxy density regions, e.g. the IR luminous fraction increases from $\\sim8$\\% at $\\Sigma\\sim10$~gal~Mpc$^{-2}$ to $\\sim25$\\% at $\\Sigma\\gtrsim100$~gal~Mpc$^{-2}$.  The observed increase is a reversal of the well-documented trend at $z<1$ and signals that we have reached the epoch when massive cluster galaxies are still forming a substantial fraction of their stars. ",
        "introduction": "A well-established observational hallmark of how galaxies evolve as a function of environment is the star formation-galaxy density relation. A plethora of studies utilizing multi-wavelength tracers of activity have shown that star formation universally decreases with increasing galaxy density at $z<1$ \\citep[e.g.][]{hashimoto:98,ellingson:01,gomez:03,patel:09}.  In particular, the cores of massive galaxy clusters are galaxy graveyards full of massive spheroidal systems that are dominated by old stellar populations.  However, as we approach the epoch when these quiescent behemoths should be forming the bulk of their stars \\citep[$z\\gtrsim2$;][]{vandokkum:98b,jorgensen:06,rettura:10}, the star formation-density relation should weaken and possibly reverse. Identifying when star formation is quenched as a function of galaxy mass and environment provides strong constraints on galaxy models \\citep[e.g.][]{kauffmann:93,hopkins:08}, i.e. is individual galaxy mass or the larger scale environment the primary driver of evolution?  Observations of field galaxies at $z\\sim1$ indicate that the star formation-density relation turns over at this epoch such that there is an enhancement of activity in the highest density regions of the field \\citep{elbaz:07,cooper:08}.  Studies also find an excess of dust-obscured star formation in group environments at $z<1$ \\citep{koyama:08,tran:09,gallazzi:09}, and very recent results suggest that star formation may be enhanced in the significantly richer core of a galaxy cluster at $z\\sim1.46$ \\citep{hilton:10}. As cluster surveys push to higher redshifts ($z>1$) and thus closer to the epoch when massive galaxies are forming their stars, variations in age will become evident in, e.g. a larger scatter in color, and robust star formation rates should reveal increasing levels of activity even in the cluster cores.  We report here the first confirmation of increasing star formation activity with increasing galaxy density observed in cluster \\irc1, a Spitzer-selected galaxy cluster at $z=1.62$ \\citep[hereafter Pap10]{papovich:10}.  We use cosmological parameters $\\Omega_m=0.3$, $\\Lambda=0.7$, and $H = 70$ km s$^{-1}$ Mpc$^{-1}$ throughout the paper; at $z=1.62$, this corresponds to an angular scale $0.5$~Mpc arcmin$^{-1}$.  All magnitudes are in the AB system. ",
        "conclusions": "We measure the rest-frame colors (dust-corrected), IR luminosities, star formation rates, and stellar masses of galaxies in \\irc, a Spitzer-selected cluster at $z=1.62$, by fitting spectral energy distributions to photometry in 10 bands ($0.4\\mu$m$<\\lambda_{obs}<8\\mu$m) and with deep \\mipsmu\\ imaging.  The cluster sample \\citep[Pap10;][]{tanaka:10} is composed of 12 spectroscopically confirmed members and 80 members selected from photometric redshifts measured using EAZY \\citep{brammer:08}; all members are within $R_{proj}\\lesssim1$~Mpc of the massive cluster galaxy located at the peak of the X-ray emission.  The 92 cluster members have a color-stellar mass distribution that is surprisingly similar to that observed in field galaxies at $z\\sim2$. When corrected for dust, the cluster members define a strong blue sequence and span a range in color, indicating a substantial amount of recent and ongoing star formation in the cluster core.  This dramatic level of activity is underscored by the 17 members detected at \\mipsmu.  In the cluster core ($R_{proj}<0.5$~Mpc), the star formation rate density from the IR luminous members alone is $\\sim1700$\\msunyr~Mpc$^{-2}$; the true value is likely to be higher given that we only include members with SFR$_{IR}\\gtrsim40$\\msunyr. These IR luminous members also follow the same trend of increasing star formation with stellar mass that is observed in the field at $z\\sim2$.  We discover the striking result that the the relative fraction of star-forming galaxies increases with increasing local galaxy density in \\irc, a reversal of the well-established trend at lower redshifts and in line with recent work at $z\\sim1.46$ that suggests enhanced star formation in cluster cores.  Measurements using star formation rates derived from the \\mipsmu\\ imaging and from the SED fitting provide independent confirmation that the relative fraction of star-forming galaxies triples from the lowest to highest density regions.  By pushing into the redshift desert ($z\\gtrsim1.6$), we are able to reach the epoch when massive cluster galaxies are still forming a significant number of their stars."
    },
    "1005/1005.0703_arXiv.txt": {
        "abstract": "The protostellar jet driven by L1448C was observed in the SiO $J$=8--7 and CO $J$=3--2 lines and 350 GHz dust continuum at $\\sim$1\\arcsec resolution with the Submillimeter Array (SMA). A narrow jet from the northern source L1448C(N) was observed in the SiO and the high-velocity CO. The jet consists of a chain of emission knots with an inter-knot spacing of $\\sim$2\\arcsec (500 AU) and a semi-periodic velocity variation. These knots are likely to be the internal bow shocks in the jet beam that were formed due to the periodic variation of the ejection velocity with a period of $\\sim$15---20 yr. The innermost pair of knots, which are significant in the SiO map but barely seen in the CO, are located at $\\sim$1\\arcsec (250 AU) from the central source, L1448C(N). Since the dynamical time scale for the innermost pair is only $\\sim$10 yr, SiO may have been formed in the protostellar wind through the gas-phase reaction, or been formed on the dust grain and directly released into the gas phase by means of shocks. It is found that the jet is extremely active with a mechanical luminosity of $\\sim$7 $L_{\\odot}$, which is comparable to the bolometric luminosity of the central source (7.5 $L_{\\odot}$). The mass accretion rate onto the protostar derived from the mass-loss rate is $\\sim$10$^{-5}$ $M_{\\odot}$ yr$^{-1}$. Such a high mass accretion rate suggests that the mass and the age of the central star are 0.03--0.09 $M_{\\odot}$ and (4--12)$\\times$10$^3$ yr, respectively, implying that the central star is in the very early stage of protostellar evolution. The low-velocity CO emission delineates two V-shaped shells with a common apex at L1448C(N). The kinematics of these shells are reproduced by the model of a wide opening angle wind. The co-existence of the highly-collimated jets and the wide-opening angle shells can be explained by the ^^ ^^ unified {\\it X}-wind model'' in which highly-collimated jet components correspond to the on-axis density enhancement of the wide-opening angle wind. The CO $J$=3--2 map also revealed the second outflow driven by the southern source L1448C(S) located at $\\sim$8.3\\arcsec (2000 AU) from L1448C(N). Although L1448C(S) is brighter than L1448C(N) in the mid-IR bands, the momentum flux of the outflow from L1448C(S) is two or three orders of magnitude smaller than that of the L1448C(N) outflow. It is likely that the evolution of L1448C(S) has been strongly affected by the powerful outflow from L1448C(N).% ",
        "introduction": "It is known that several deeply-embedded young stellar objects are driving highly-collimated molecular outflows with an extremely-high velocity (EHV) and jet-like component running along the axes of the lobes \\citep[e.g.,][]{Bac96}. Since the EHV molecular jets having the terminal velocities of 50--150 km s$^{-1}$ are concentrated within the narrow angles along the axes of the lobes and have large momenta comparable to those of the slowly moving (20--30 km s$^{-1}$) ^^ ^^ classical'' outflows, the EHV jets are considered to be closely connected to the ^^ ^^ primary jet'' which is responsible for driving molecular outflows. Therefore, studying the physical and kinematical properties of the EHV jet will allow us to understand the properties of the ^^ ^^ primary jet'' which provide clues to constrain the launching mechanism of the outflows.  The low-luminosity (7.5 $L_{\\odot}$; \\citet{Tob07}) class 0 source L1448C (also known as L1448-mm) in the Perseus molecular cloud complex (D$\\sim$250 pc: e.g., \\citet{Eno06}) is a spectacular example of an outflow with an EHV jet. The EHV component of this source was identified as the secondary peaks of the CO $J$=2--1 spectra at ${\\sim}{\\pm}$60 km s$^{-1}$ from the cloud systemic velocity \\citep{Bac90}. The CO emission in the EHV range was found to be confined to a series of discrete clumps, called ^^ ^^ bullets\", which are aligned along the axes of the outflow lobes and are symmetrically placed with respect to the central source. The EHV bullets are also observed in several transitions of the SiO \\citep[e.g.][]{Bac91, Dut97}. Since the SiO emission has been barely detected in quiescent dark clouds because of its very low abundance (of the order of 10$^{-12}$; \\citet{Ziu89,Mar92}), the detection of SiO in EHV bullets suggests the presence of shocks that enhanced the SiO abundance in bullet gas by a factor of $\\gtrsim$10$^4$. Although the lower transition of SiO, i.e. $J$=2--1,  was observed not only in the EHV bullets but also in the lower velocity component that delineates the tips and walls of the outflow cavities, $J$=5--4 emission was confined to a pair of EHV bullets located at the closest positions to the star \\citep{Bac91}. This suggests that the excitation condition of the EHV jet varies along the axes, and that the jet gas is highly excited in the close vicinity of the driving source. Recently, higher excitation SiO up to $J$=11-10 has been observed by \\citet{Nis07}. Their results have revealed that the innermost pair of bullets, labeled B1 and R1 by \\citet{Bac90}, have a density of $n_{\\rm H_2}\\sim10^6$ cm$^{-3}$ and a kinetic temperature of $T_{\\rm kin}\\gtrsim$ 500 K, which is denser and warmer than the bullets in the downstream. It is also known that the innermost pair of bullets, B1 and R1 are resolved into two clumps, BI-BII and RI-RII, respectively, in the higher resolution ($\\sim$2\\arcsec) interferometric SiO $J$=2--1 observations \\citep{Gui92, Gir01}. The high resolution SiO data exhibit the kinematic structure of the EHV jet near the source, which shows an apparent acceleration of the jet up to 70 km s$^{-1}$ within a region of 6$''$ ($\\sim$2000 AU). The proper motion measurement of the SiO clumps carried out by \\citet{Gir01} suggests that the outflow axis is inclined by 21$^{\\circ}$ with respect to the plane of the sky, and therefore, the SiO clumps in the EHV jet are likely to be moving with absolute velocities of 180 km s$^{-1}$.   In this paper, we present the SiO $J$=8--7, CO $J$=3--2, and 350 GHz continuum images obtained with the Submillimeter Array\\footnote{The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica.} (SMA) \\citep{Ho04} at $\\sim$1 arcsecond resolution, which is a factor of three higher than the previous SiO $J$=2--1 images \\citep{Gui92, Gir01}. A high angular resolution is crucial for studying the structure and kinematics of the jet near the base, at which the jet velocity increases up to the terminal velocity. In addition, higher transitions of SiO and CO in the submillimeter waverange enable us to segregate the dense and warm gas in the EHV jet from the lower excitation gas in the cavity wall. ",
        "conclusions": "The central region of the highly-colimated molecular outflow driven by L1448C was mapped in the SiO $J$=8--7, CO $J$=3--2, and 350 GHz continuum emission with the SMA at $\\sim$1 arcsecond resolution. Our main conclusions are the following: \\begin{enumerate} \\item The 350 GHz continuum emission was detected from two {\\it Spitzer} sources L1448C(N) and L1448C(S). The continuum emission from L1448C(N) consists of an extended component  and a compact component. The compact component is elongated perpendicular to the outflow axis, and is likely to be a circumstellar disk with a size of $\\sim$90 AU. The spectral index of this compact component derived from the data from 86 GHz to 350 GHz is ${\\alpha}{\\sim}$2, suggesting the possibility that the continuum emission is optically thick at 350 GHz. The mass of the disk is estimated to be $\\sim$0.11 $M_{\\odot}$. \\item The continuum flux from L1448C(S) is $\\sim$60 mJy, which is $\\sim$10 times lower than the flux from L1448C(N), although L1448C(S) is brighter than L1448C(N) in the mid infrared wavebands. The mass of the circumstellar material surrounding L1448C(S) is estimated to be 8.6$\\times$10$^{-3}$ $M_{\\odot}$. \\item A narrow jet from L1448C(N) along the outflow axis was observed in the SiO and the high-velocity CO. The width of the jet measured in the SiO images is $\\sim$200 AU FWHM on average. The jet consists of a chain of emission knots with an inter-knot spacing of $\\sim$500 AU. It is likely that the knots in the L1448 jet are the internal bow shocks in the jet beam. \\item The dynamical timescale of the innermost pair of knots, which are significant in the SiO but barely seen in the CO, is only $\\sim$10 yr. It is likely that the SiO may have been formed quickly in the protostellar wind through the gas-phase reaction, or been formed on the dust grain and directly released into the gas phase by means of shocks. \\item The low velocity CO emission delineates two V-shaped shells with a common apex at  L1448C(N). The kinematics of this shell component is reproduced by the model of wide opening angle wind. Therefore, the outflow from L1448C(N) consists of both highly-collimated jets and shells produced by wide-opening angle wind. The observed jet+shell structure can be explained by the ^^ ^^ unified model'' proposed by \\citet{Sha06}, in which highly-collimated jet components are explained as an on-axis density enhancement of the {\\it X}-wind type of wide opening angle wind. \\item The Jet from L1448C(N) is extremely active with a momentum supply rate of $\\sim$5$\\times$10$^{-4}$ $M_{\\odot}$ km s$^{-1}$ yr$^{-1}$ and a mechanical luminosity of $\\sim$7 $L_{\\odot}$. The mass accretion rate derived from the mass loss rate  is $\\sim$10$^{-5}$ $M_{\\odot}$ yr$^{-1}$. Such a high mass-accretion rate and a rather low bolometric luminosity of the central source, 7.5 $L_{\\odot}$, imply that the central protostar is still in the very early phase of its evolution with a mass of 0.03--0.09 $M_{\\odot}$ and a dynamical age of (4--12)$\\times$10$^3$ yr. \\item The blue part and the red part of the jet are misaligned by $\\sim$5$^{\\circ}$ and forming a C-shaped bending toward the west. The possible origin of this bending is the dynamical pressure caused by the outflow from the nearby protostar, L1448N. \\item The jet is deflected toward the east in the blueshifted part and toward the south in the redshifted part. In addition, the jet is wiggling with a period of $\\sim$15--20 yr. The deflection and wiggling of the jet can be explained if the driving source is a member of the binary system with an orbital radius of 2--4 AU. \\item The jet shows a semi-periodic variation in radial velocity with an amplitude of $\\sim$7 km s$^{-1}$. Each SiO knot has its higher velocity in the upstream side and lower velocity in the downstream side. The opposite velocity gradient is seen in the faint emission between the knots. It is likely that the ejection velocity varies periodically by means of modulation of mass accretion. \\item The bipolar outflow in the NE-SW direction centered at L1448C(S) was discovered in the CO $J$=3--2. This provides a strong evidence that L1448C(S) is a protostar. The momentum flux of this outflow is only $\\sim$10$^{-6}$ $M_{\\odot}$ km s$^{-1}$ yr$^{-1}$, which is two to three orders of magnitude smaller than that of L1448C(N) outflow, and is comparable to those of class I outflow. \\item L1448C(S) is surrounded by a rather small amount of circumstellar material of less than 0.01 $M_{\\odot}$ and is powering a less energetic outflow, suggesting that this source has a nature close to class I rather than class 0, even though this source is formed in the same dense core as L1448C(N). One possible scenario to explain this dichotomy is the effect of the outflow from L1448C(N); significant amount of material in the envelope surrounding L1448C(S) might have been stripped off by the powerful outflow from L1448C(N). Another possibility is the disintegration of an unstable multiple system, in which L1448C(S) is an escaping member with a truncated disk.  \\end{enumerate}"
    },
    "1005/1005.3700_arXiv.txt": {
        "abstract": "We study the luminosity function (LF), the comoving rate and the detection rate of Long Gamma-Ray Burst (LGRBs) to high redshift, using galaxy catalogues constructed by combining high-resolution N-body simulations with semi-analytic models of galaxy formation. We assume the collapsar model and different metallicity thresholds, and conclude that LGRBs are not good tracers of the star formation history in the universe. Then using the $\\log N-\\log P$ diagram for BATSE bursts, we determine the LF (with and without evolution with redshift) and the formation rate of LGRBs, obtaining constraints on the slope of the power-law. We check the resulting redshift distribution with \\sw\\ data updated to 2009 August, finding that models where LGRBs have as progenitors stars with $Z<0.3Z_{\\odot}$ and without evolution of the LF are in agreement with the data. We also predict that there are about $\\sim1\\%$ of GRBs at redshift $z>6$. ",
        "introduction": "\\label{sec:intro} Gamma-ray bursts (GRBs) are the most luminous explosion in the Universe \\citep{zhang04}. They offer exciting possibilities for studying astrophysics in extreme conditions, e.g., radiative processes in highly relativistic ejecta \\citep[and references therein]{huang00,fan08}. Because of their very large luminosity, GRBs represent cosmological events, which have been detected up to $z\\sim8.2$ \\citep{tan09,sal09b}.  The observed distribution of the duration of GRBs is bimodal \\citep{kou93}: long GRBs (hereafter LGRBs) and short GRBs, depending on whether their durations are longer or shorter than 2 seconds. The current favorite hypotheses for their origin are that short GRBs are produced by the merger of compact objects \\citep{li98,osh08}, while LGRBs originate from the death of massive stars (with low metallicity), such as Wolf-Rayet stars ({\\sl the collapsar model}) \\citep{Yoon_Langer_Norman_2006,yoon08, woo06b}. Throughout this paper we will deal only with the LGRBs. Observational data are consistent with the hypothesis of the LGRB-supernova connection: at least some LGRBs are associated with core-collapse supernovae \\citep[and references therein]{gal98,hjo03,pia06,li06,woo06b}. In addition, all supernovae associated with GRBs are Type Ic, which supports the hypothesis of Wolf-Rayet stars as progenitors of LGRBs. Because of their connection with supernovae, LGRBs are potential tracers of the cosmic star formation history (SFH) \\citep{bro02,fyn06,pri06,pro07,Savaglio_2006,tot06,li08a}, nonetheless this connection could be non-trivial (see e.g. \\cite{koc09}).  The detection of LGRBs at high redshift and the connection with Type Ic SNe, make them promising for probing the Universe. They are probably the only objects that allow us to study the cosmos {at high redshift} and the early evolution of {PopIII stars}. Therefore, understanding where {GRBs} are distributed in the universe and how they are connected with the star formation rate (SFR) is very {important}. Since the launch of the \\sw\\ satellite \\citep{geh04}, the number of GRBs with measured redshift has been greatly increased. Nonetheless to date there are only $\\sim140$ \\sw\\ GRBs with known redshifts. This sample is still too {small} to constrain their luminosity function (hereafter LF).  Previous studies \\citep[and others]{Porciani_Madau_2001, gue05, nat05, dai06, sal07} {touched} the problem of determining the redshift distribution of GRBs differently. Their method is based on constraining the GRBs distribution by assuming an average energy spectrum for all the bursts and that {GRBs} trace the evolving SFR, either with a constant or evolving LF. The redshift distribution, together with the LF, can provide important insights not only into the physics of the individual objects themselves, but also into the evolution of matter in the Universe.  One of the first pioneering {works} was Porciani $\\& $ Madau (2001). The rate of GRB is {fitted to} observational data, using the assumption that SFR is proportional to the GRB rate \\citep{hop06}. Finally {the rate is} convolved with a selection function dependent on the instrument used. They find that the rate of {bursts} is of about 1-2 GRBs for every one million Type II SNe. Notable recent attempts include: Guetta, Piran $\\& $ Waxman (2005) explore a variety of different star-formation rate histories and GRB luminosity {functions}; Natarajan et al. (2005) additionally incorporate a simple prescription for a low-metallicity preference {of GRBs}. Daigne et al. 2006 used Monte Carlo simulations to predict {the GRB evolution}, assuming that GRBs follow the SFR, the LF is a power-law ({independent of} redshift) and the peak energy is determinated by {two relations}. They find that: the slope of the LF is between $1.5-1.7$, the Amati relation should be {an intrinsic relation} and, the GRB rate density at $z=7$ is about 6-7 times larger than at $z=2$. They {also deduced} that the properties of GRBs and GRB-progenitors are redshift dependent, since the redshift distribution of \\sw\\ burst strongly favors their SFR3 model (see Daigne et al 2006), {although that is} an unrealistic model.\\\\ Salvaterra follows this approach in two different works (Salvaterra $\\&$ Chincarini 2007, Salvaterra et al. 2008), assuming that GRB luminosity evolves with redshift and that GRBs form preferentially in low-metallicity environments. They use this constraint to set a robust upper limit on the {slope of} bright-end of GRB LF, finding that the number of bright {GRBs} detected by \\sw\\ implies that {it} cannot be very steep ($\\delta<2.6$ for progenitors with $Z<0.3Z_{\\odot}$). Moreover they found that assuming a threshold of $F>0.4$ [ph/cm$^2$/s], at least $\\sim 5-10\\%$ of all detected GRB should lie at redshift $z>5$.\\\\  In this work, we do not use a GRB comoving rate proportional to the star formation in the Universe, but only we assume that {the global rate of GRBs} per SNe is on average (over all cosmic times) of about 1 GRB event every 1000 SNe \\citep{lan06}. We derive the LF and formation rate of GRBs using a catalogue of {galaxies} constructed by combining high-resolution N-body simulations with a semi-analytic model of galaxy formation (Wang et al. 2008). We fit the observed logN-logP relation \\citep{kom00} derived from {the GRB data of} the Burst And Transient Source Experiment on board the {\\it CGRO} satellite (BATSE, Fishman et al. 1989) in order to constrain the free parameters of the LF. We adopt this method for three {GRB progenitor subsamples} with different cuts in metallicity, following the collapsar model, and assuming a constant and evolving LF.\\\\ By comparing the cumulative distribution of peak photon fluxes, \\citet{dai09} recently proved that the \\sw\\ and BATSE samples track the same parent population of bursts. For this reason, following Porciani $\\& $ Madau, we rely on the {GRBs} observed by BATSE as {the two} samples should have comparable LF.  The paper is organized as follows. We present in section 2 {the simulated GRB sample} used in this work. In section 3, we describe the method to reproduce the redshift distribution of LGRB. We describe our results in section 4. We give our conclusions in section 5. ",
        "conclusions": "In this work, we have studied the luminosity function (LF), the comoving rate and the detection rate of Long Gamma-Ray Burst (LGRBs) in the context of a hierarchical model of galaxy formation.  Assuming the collapsar model and imposing different metallicity constraints we find that:\\\\ \\begin{itemize} \\item GRBs with low metallicity progenitors ($Z<0.1-0.3Z_{\\odot}$) do not represent a perfect tracer of the cosmic star formation history (see Fig.1). The deviation of the LGRB rate from the star formation rate decreases with increasing redshift (as a consequence of the global decrease of metallicity with increasing redshift) and the bias is stronger as the metallicity threshold assumed is lowered.\\\\  \\item The LF of LGRBs is well descripted by a power-law with exponential cut-off, with well determined slope between $1.686<\\xi<1.838$. Conversely, the characteristic luminosity spans a large range of values ($0.038<L_0<77.875\\times 10^{50}$[erg/s]), increasing at lower metallicity threshold (see table 1).\\\\  \\item It is possible to reproduce, within error bars, the redshift distribution of a subset\\footnote{Only GRBs well placed for optical observations, \\cite{jak06}} of \\sw\\ data with $F>0.2$ [ph/cm$^2$/s], using: (I) a model without cut in metallicity with a very strong evolving LF, or (II) a model with metallicity threshold $Z<0.3Z_{\\odot}$ and a non-evolving LF (see Fig.3). Selecting only {Bright} LGRBs ($F>1\\,$[ph/cm$^2$/s]) and suspecting that a scenario where the properties of GRB change so strong with redshift seem to be enough unrealistic, we rule out the (I) model (see Fig. 4).\\\\  \\item We predict to have $\\sim1\\%$ Bright bursts {at} high redshift ($z>6$).\\\\  \\end{itemize}   Our work constrains the LF function of LGRBs, using a rate of GRB not proportional to SFR, and assuming different metallicity for the progenitors, giving us the possibility to assert that the most probable model for LGRB's progenitor have $Z<0.3Z_{\\odot}$ and {no} evolution in the LF."
    },
    "1005/1005.1643_arXiv.txt": {
        "abstract": "Cosmological simulations of dark matter structures have identified a set of universal profiles, and similar characteristics have been seen in non-cosmological simulations. It has therefore been speculated whether these profiles of collisionless systems relate to accretion and merger history, or if there is an attractor for the dark matter systems.  Here we identify such a 1-dimensional attractor in the 3-dimensional space spanned by the 2 radial slopes of the density and velocity dispersion, and the velocity anisotropy. This attractor effectively removes one degree of freedom from the Jeans equation. It also allows us to speculate on a new fluid interpretation for the Jeans equation, with an effective polytropic index for the dark matter particles between $1/2$ and $3/4$. If this attractor solution holds for other collisionless structures, then it may hold the key to break the mass-anisotropy degeneracy, which presently prevents us from measuring the mass profiles in dwarf galaxies uniquely. ",
        "introduction": "Cosmological dark matter structures which have been formed in an expanding universe appear to have virtually universal profiles.  This was first emphasized for the density profile through numerical simulation when \\cite{nfw} demonstrated that all structures from galaxies to galaxy clusters are well fitted by the same simple form \\begin{equation} \\rho = \\frac{\\rho _0}{\\frac{r}{r_s} \\left( 1  + \\frac{r}{r_s} \\right) ^2} \\, , \\label{eq:nfw} \\end{equation} where $r_s$ is the scale radius. Even though various theoretical attempts have been made to explain such density profile (e.g. \\cite{henriksen,gsmh,hekang,hjortwilliams}) there is still no concensus of the underlying physical reason for this universality.   To ease the task of theoreticians one should look for simple connections in some para\\-meter space. One such connection was suggested by \\cite{taylornavarro}, who presented numerical evidence that the pseudo phase-space density is a power-law in radius \\begin{equation} \\frac{\\rho}{\\sigma^3} \\sim r^{-\\alpha} \\, , \\label{eq:nt} \\end{equation} where $\\alpha = 1.875$ is in agreement with the predictions from the spherical infall model \\citep{bertschinger}.  This numerical evidence was soon used in conjunction with the Jeans equation to derive density profiles in general agreement with eq.~(\\ref{eq:nfw}) \\citep{hansenjeans, austin, dehnenmclaughlin}.  One problem does, however, present itself, namely that recent high resolution simulations do not confirm that the pseudo phase-space density is a universal straight line (for recent discussions, see \\cite{schmidt09, knollmann, ludlow}).   In this Letter we study the evolution of pure dark matter structures using numerical simulations. We find that all the various equilibrium dark matter structures, when perturbed in a simple but realistic manner, move towards a particular curve in the 3-dimensional space of the derivatives of density and velocity dispersion, and the velocity anisotropy. ",
        "conclusions": "We have identified an attractor solution for dark matter structures. This implies that any dark matter structure which is repeatedly perturbed (e.g. through violent relaxtion during merging) and then allowed to relax (phase mix), will flow towards this 1-dimensional curve in the 3-dimensional space spanned by the 2 radial derivatives of the density and velocity dispersion, and the velocity anisotropy.   This finding provides strong support for the idea that the universalities found in cosmological dark matter structures are a property of gravity, and not simply a result of similar accretion and merger histories of different structures.  This attractor solution effectively removes one degree of freedom from the Jeans equation, giving hope that we will eventually be able to solve the Jeans equations analytically, and thereby truely understand the origin of the universal profiles.  \\noindent  {\\bf Acknowledgements}\\\\ It is a pleasure to thank Jens Hjorth for discussions. The simulations were performed on the facilities provided by the Danish Center for Scientific Computing. The Dark Cosmology Centre is funded by the Danish National Research Foundation."
    },
    "1005/1005.4066_arXiv.txt": {
        "abstract": "Cosmological hydrodynamic simulations robustly predict that high-redshift galaxy star formation histories (SFHs) are smoothly-rising and vary with mass only by a scale factor.  We use our latest simulations to test whether this scenario can account for recent observations at $z\\geq6$ from WFC3/IR, NICMOS, and IRAC.  Our simulations broadly reproduce the observed ultraviolet (UV) luminosity functions and stellar mass densities and their evolution at $z=6$--8, all of which are nontrivial tests of the mean SFH.  In agreement with observations, simulated galaxies possess blue UV continua owing to young ages (50--150 Myr), low metallicities (0.1--0.5$Z_\\odot$), and low dust columns ($E(B-V) \\leq 0.05$).  Our predicted Balmer breaks at $z=7$, while significant, are $\\approx0.5$ magnitudes weaker than observed even after accounting for nebular line emission, suggesting observational systematic errors and/or numerical resolution limitations.  Observations imply a near-unity slope in the stellar mass--star formation rate relation at all $z=6$--8, confirming the prediction that SFH shapes are invariant.  Dust extinction suppresses the UV luminosity density by a factor of 2--3, with suppression increasing modestly to later times owing to increasing metallicities.  Current surveys detect the majority of galaxies with stellar masses exceeding $10^9\\msun$ and few galaxies less massive than $10^{8.5}\\msun$, implying that they probe no more than the brightest $\\approx30\\%$ of the complete star formation and stellar mass densities at $z\\geq6$.  Finally, we demonstrate that there is no conflict between smoothly-rising SFHs and recent clustering observations.  This is because momentum-driven outflows suppress star formation in low-mass halos such that the fraction of halos hosting observable galaxies (the ``occupancy\") is 0.2--0.4 even though the star formation duty cycle is unity.  This leads to many interesting predictions at $z\\geq4$, among them that (1) optically-selected and UV-selected samples largely overlap; (2) few galaxies exhibit significantly suppressed specific star formation rates; and (3) occupancy is constant or increasing with decreasing luminosity.  These predictions are in tentative agreement with current observations, but further analysis of existing and upcoming data sets is required in order to test them more thoroughly. ",
        "introduction": "\\label{sec:intro} In the era of precision cosmology, the study of galaxy evolution is an initial value problem in which the goal is to characterize the processes that connect the initial conditions at $z\\approx1090$~\\citep[for example,][]{kom10} to galaxies, which had evidently begun forming by $z=10$~\\citep[for example,][]{sta07,bou09c,yan09}.  The initial-value nature of the problem implies that detailed measurement of its earliest stages constitutes an immensely useful aid through its ability to anchor an understanding of the later stages.  Gas inflows, outflows, and other feedback processes impact high-redshift star formation histories (SFHs).  Hence observational inferences regarding the SFHs at high redshift constrain star formation and feedback both at early times and at subsequent epochs.  Cosmological hydrodynamic simulations robustly predict smoothly-rising SFHs~\\citep{fin07} with a scale-invariant shape at $z\\geq6$.  Therefore, any comparison between simulation predictions and observations constitutes a (possibly indirect) test of this scenario.  We have previously demonstrated considerable success in using cosmological hydrodynamic simulations to interpret observations of high-redshift galaxies.  In~\\citet{fin06}, we showed that simulations of a \\lcdm~universe with star formation modulated by outflows with a constant wind speed and mass loading factor broadly reproduce the observed rest-frame UV luminosity function (LF) at $z\\sim4$.  We confirmed previous numerical predictions~\\citep[for example,][]{dav00,wei02} that stellar mass ($\\mstar$) and star formation rate (SFR) correlate, finding a small scatter and a slope near unity.  A qualititatively similar correlation has now been observed at all redshifts out to $z=7$~\\citep{bri04,elb07,noe07a,dad07,sal07,sch07,pan09,lab09,bot09,oli10,mag10}, although observations at $z<2$ tend to suggest a somewhat flatter slope of 0.7--0.9 that is associated with galaxy downsizing~\\citep{bri04,elb07,noe07a,dad07,sal07,sch07,bot09,mag10}.  This correlation has two implications: First, it implies that UV-selected galaxies are generically the most massive (star-forming) galaxies at a given redshift~\\citep{nag04,fin06} unless the scatter is quite large~\\citep{wei02}.  Second, a slope of unity supports a scenario in which SFHs have a scale-invariant shape such that, on average, galaxies' SFHs differ only by a scale factor.  The SFR-$\\mstar$ relation therefore constitutes a key prediction of galaxy evolution models~\\citep[for example,][]{bou09,dut09}.  We additionally predicted that high-redshift galaxies should exhibit pronounced Balmer breaks because smooth SFHs naturally yield evolved stellar populations.  With the arrival of rest-frame optical constraints from IRAC, this prediction has now received dramatic confirmation~\\citep[for example,][]{bra08,cha05,dow05, dun07,ega05,eyl05,eyl07,lab06,lai07,lab09,sta09,ver07,yan05,pen09,zhe09}, although the question of whether emission lines could be mimicking the observed Balmer breaks remains unresolved~\\citep{sch09,sch10}.  In~\\citet[hereafter DFO06]{dav06}, we tested a treatment for momentum-driven outflows from star-forming regions against the available constraints on the rest-frame UV and \\lya LFs at $z\\sim6$, finding reasonable agreement in both cases.  This result was subsequently extended by~\\citet{bou07}, who found that our simulation's prediction for the evolution of the UV LF normalization from $z=8\\rightarrow4$ was well-matched by observations.  This is a nontrivial test of our predicted SFHs because constant or decaying SFHs would predict a different evolution of the UV LF than is observed.  Finally, in~\\citet{fin07}, we showed that our simulations could account for the rest-frame UV-to-optical spectral energy distributions (SEDs) of 5 out of 6 observed galaxies between $z=5.5$--6.5.  We demonstrated that smoothly-rising SFHs are equally as plausible as more popular constant or exponentially-decaying models.  We then used these comparisons to explore the implied physical properties such as SFR, $\\mstar$, and age, finding good agreement with previous constraints although with significantly narrower posterior probability distributions owing to the tight priors imposed by the physical correlations that our simulations predict.  While these studies are encouraging, their support for our SFH scenario is somewhat indirect because, until recently, observations at $z\\geq6$ did not clearly imply a unique SFH.  For example, any galaxy evolution model that predicted scale-invariant SFH shapes would predict an SFR-$\\mstar$ relation with a slope of unity.  Similarly, any SFH shape can be tuned to reproduce the observed UV LF at one epoch, and bursty SFHs could be tuned to reproduce it at multiple epochs.  The fundamental difficulty is that broadband SEDs do not contain enough information to constrain the underlying SFHs of individual objects, even when analyzed through detailed SED-fitting studies~\\citep[for example,][]{sha01,pap01}.  Since the completion of these previous works, a number of authors have used new data from the HST WFC3/IR and NICMOS cameras as well as Spitzer/IRAC to increase the number of objects at $z\\geq6$ with published rest-frame UV to optical constraints \\citep{bou09b,bou10a,bun09,cas09,fink09,gon09,lab09,lab10,mcl10, oes10a,oes10b,ouc09,yan09}.  Using these new samples, we may now ask whether the \\emph{mean} SFH of reionization-epoch galaxies is constrained even if \\emph{individual} SFHs are not. For example,~\\citet{sta09} recently pointed out that the observed lack of evolution in the specific star formation rate (SSFR) of UV-selected galaxies from $z=$4--6~\\citep[see also][]{ove09,gon09} rules out the possibility that the majority of galaxies observed at $z=4$ have been forming stars at constant or declining rates since $z\\geq5$.  This is because constant or declining SFHs would predict higher SSFRs at earlier times, in conflict with their observations.  They noted that smoothly-rising SFHs could explain their observations since in this case galaxies would grow in $\\mstar$ and SFR simultaneously.~~\\citet{pap10} studied the evolution of the UV luminosity and stellar mass at constant number density and concluded that observations require smoothly-rising SFHs and are consistent with a scale-invariant shape from $z=8\\rightarrow 3$. Finally,~\\citet{mar10} argued in favor of exponentially-rising over decaying SFHs as models for star-forming galaxies at $z\\sim2$ for two reasons.  First, they found that rising SFHs provide better fits both to observations and to predictions from semi-analytic models of galaxy formation.  Second, rising SFHs naturally explain higher-redshift samples as the fainter progenitors of lower-redshift samples.  These suggestions exemplify how statistical samples at multiple epochs can constrain the mean SFH even if the SFHs of individual galaxies are unconstrained.  In this work, we use this idea to build upon the results of~\\citet{fin07} by asking whether smoothly-rising SFHs can explain the observed statistical properties of galaxies at $z=6$--8. Broadly, since all of our simulated high-redshift galaxies experience smoothly-rising SFHs, any comparison with observations can be viewed as an indirect test of these models.  In most cases, however, the constraining power of large samples at multiple epochs will render our comparisons direct tests of the smoothly-rising scenario.  We describe our simulations in Section~\\ref{sec:sims}.  In Section~\\ref{sec:meanSFH}, we review the rising SFH scenario and summarize its predictions.  We compare our predicted UV LFs with observations in Section~\\ref{sec:lfs}.  We perform detailed comparisons between simulated and observed SEDs in Section~\\ref{sec:colors}. Additionally, we use our simulations to predict the completeness of current $z\\geq6$ surveys.  In Section~\\ref{sec:phys}, we study the predicted relationships between stellar mass, star formation rate, star formation history, and metallicity, comparing with observations where possible.  In Section~\\ref{sec:occ}, we use our predicted halo occupation distribution to argue that recent clustering observations are consistent with smooth SFHs, and we discuss further tests of our model and alternative bursty SFH scenarios.  In Section~\\ref{sec:lowz}, we comment on processes that cause SFHs to depart from our smoothly-rising, scale-invariant scenario at lower redshift. Finally, we summarize our results in Sec~\\ref{sec:sum}. ",
        "conclusions": "\\label{sec:sum} We have conducted a detailed comparison between the predictions of state-of-the-art cosmological hydrodynamic simulations including momentum-driven outflows and recent observations of reionization-epoch galaxies in order test two fundamental predictions regarding high-redshift SFHs: \\begin{enumerate} \\item SFHs at early times are smoothly-rising; and \\item SFH shapes are scale-invariant. \\end{enumerate} Our findings are as follows: \\begin{enumerate} \\item Our simulations reproduce the observed rest-frame UV luminosity function as well as its integral to current limits to within a factor of 2--3 during the epoch $z=$6--8.  They also reproduce the observed stellar mass density to current limits.  Either of these constraints could in principle be matched by any SFH at a single epoch, but the only smooth SFH that can account for their evolution is a rising one.  This suggests that our prescriptions for star formation and feedback yield SFHs in reasonable agreement with current constraints.  In detail, we find that increasing the assumed outflow mass-loading factors by $\\approx$50\\% could bring the predicted and observed UV luminosity densities into better agreement without compromising the good agreement with the observed stellar mass density (to within the reported errors).  Integrating below current detection limits, we estimate that current observations probe the brightest 30\\% of the total star formation and stellar mass densities at $z\\geq6$. \\item We reproduce the observed blue UV continua at $z=7$ for all but the faintest galaxies and marginally reproduce the observed bluening to lower luminosities.  The UV continua of the faintest galaxies may suggest lower metallicities and a nonzero ionizing escape fraction~\\citep{bou10a}.  We attribute the observed colour-magnitude trend to a modest level of dust reddening that increases with luminosity as seen at lower redshifts~\\citep{meu99,sha01}.  This suggests that exotic stellar populations are not required in order to explain observations of galaxies brighter than $H_{160}=27.5$ at $z=7$~\\citep{fink09,bou10a}. \\item Our simulations' relatively evolved stellar populations predict red $H_{160}-[3.6]$ colours at $z\\geq6$, augmented by dust and optical emission lines.  However, the predicted  colours are $\\approx0.5$ mag bluer than reported by~\\citet{lab09}.  It is not clear whether resolving this discrepancy requires significantly higher mass resolution or improved observations.  For the brightest galaxies, allowing the redshift and extinction to vary freely enables our models to reproduce the observed $K$--$[4.5]$ fluxes while underproducing the observed $J_{125}$ and $H_{160}$ fluxes.  This suggests that the red observed $H_{160}-[3.6]$ colours may be as likely to owe to systematic error in the WFC3 as in the IRAC observations. \\item We reproduce the near-unity slope of the observed SFR-$\\mstar$ relation at all redshifts.  This agreement supports a scenario in which observed reionization-epoch galaxies began forming stars at similar epochs and possess SFHs that vary only by a scale factor, which is a robust prediction of hydrodynamic simulations.  We also reproduce the observed slow evolution of the SSFR, which is incompatible with constant or smoothly declining SFHs while supporting smoothly-rising SFHs~\\citep{sta09}.  The predicted scatter is $\\approx 0.1$dex, less than the (still remarkably tight) reported scatter of 0.25 dex.  The larger observed scatter may imply a need for higher mass resolution in order to resolve the impact of minor mergers, but it may also be boosted by observational uncertainties. An offset of a factor of 2--3 between observations and predictions suggests that observationally-inferred stellar masses should be corrected for the likely impact of optical emission lines. \\item Our significantly improved treatment for the production and transport of metals predicts that observable galaxies at $z\\geq6$ possess metallicities in excess of 0.1$Z_\\odot$.  This is several times larger than our previous prediction.  Increasing the assumed outflow mass-loading factors to improve the agreement between the predicted and observed UV LFs would lower these predictions by $\\sim0.1$ dex, but it would not change the fundamental result that galaxies self-enrich quite quickly.  Hence we confirm that it should be possible to understand current samples without reference to Population III stars.  Nevertheless, a residual fraction $10^{-3}$ of all star formation at $z=6$--8 may occur below the Population III metallicity threshold, which may eventually be visible as pair-production supernovae or GRBs. \\item There is no conflict between observations suggesting that many massive halos do not contribute to the observed samples at $z=4$--5 and the smoothly-rising SFH scenario.  This is because clustering observations constrain the halo occupancy rather than the star formation duty cycle.  The latter is only one of several possible interpretations of an overall occupancy that is less than unity, and our simulations disfavor it.  Instead, we favor the interpretation that it reflects the scatter in the luminosity-halo mass relation. This is because strong outflows boost the scatter in the SFRs of low-mass halos, suppressing the UV luminosities of many isolated central galaxies that would be visible in the absence of outflows. Consequently, more than half of observed galaxies live in halos whose mean (observable) occupancy is less than one and the overall fraction of halos above $M_{\\mathrm{min}}$ that host an observable galaxy is 0.2--0.3 as observed even though the predicted star formation duty cycle is unity.  A halo occupancy in this range should thus be regarded as indicative of a high star formation duty cycle rather than a low one.  This interpretation makes a number of predictions: \\begin{enumerate} \\item The minimum mass for hosting an observed galaxy $M_{\\mathrm{min}}$ is at least twice as large as would be expected given galaxies' baryonic masses and the global baryon mass fraction; \\item The minimum mass for hosting an average of two galaxies is larger than $M_{\\mathrm{min}}$ by a factor of 17--18; \\item Optically-selected and UV-selected samples should largely overlap at $z\\geq4$; \\item The scatter in the SFR-$\\mstar$ relationship should be small, and few galaxies should be observed with suppressed SSFRs indicative of a transition between active and quiescent states; \\item The overall occupancy for faint samples should be greater than or equal to the overall occupancy for brighter samples, in direct conflict with a picture in which lower-mass galaxies possess burstier SFHs in analogy with local dwarf galaxies. \\end{enumerate} Each of these predictions is in tentative agreement with current observations, but we expect that more detailed analyses of current and upcoming samples will test them more thoroughly.  \\end{enumerate}  Further progress in understanding high-redshift SFHs will occur on both theoretical and observational fronts.  On the theoretical side, it would be interesting to increase our mass resolution significantly while relaxing the assumption that wind particles temporarily decouple hydrodynamically in order to resolve feedback on small scales. Simulations of dwarf galaxy formation at low redshift tend to find burstier SFHs as the mass resolution is increased~\\citep[for example,][]{val08} owing to stochastic processes~\\citep{ger80}, hence it is natural to suppose that similarly low-mass objects at $z\\geq6$ would also be bursty.  On the other hand, dwarf galaxy SFRs are generally $<0.5 \\msun\\mbox{yr}^{-1}$ whereas even the faintest $z=7$ galaxies indicate SFRs of $>2\\msun\\mbox{yr}^{-1}$; whether bursty SFRs are expected in the presence of efficient gas inflows is unclear.  If so, this may improve the agreement with the observed scatter in the SFR-$\\mstar$ relation as well as the colours of low-mass galaxies~\\citep{lab09}.  Unfortunately, increasing our mass resolution is not trivial because our present simulations already probe current computational limitations.  Further progress will require resimulations of a statistical sample of high-redshift galaxies~\\citep[for example,][]{bro07}. By contrast, incorporating stronger outflows in order to improve the agreement with the observed UV luminosity function would be trivial (aside from the computational expense), although it remains to be seen whether this would also preserve the agreement with the observed density of metals in the intergalactic medium~\\citep{opp09b}.  The impact of an inhomogeneous ionizing background warrants closer scrutiny.  Our current work does not indicate that our assumption of a uniform~\\citet{haa01} ionizing background produces large errors in our predictions (DFO06).  Nonetheless, it would be useful to study the relative roles of photoionization heating versus outflows in suppressing star formation in faint galaxies using simulations that treat the nascent ionizing background self-consistently.  Preliminary work suggests that interesting nonlinear interactions exist between outflows and the ionizing background; these warrant more detailed study~\\citep{paw09}.  A closer inquiry into the nature of our predicted HOD and its ability to reproduce clustering constraints at $z=4$--5 is in order.  We have not delved into a full clustering analysis because our goal has been to demonstrate that clustering properties in the presence of mass-dependent outflows are qualitatively different than the conventional central-satellite HOD picture and have the potential to reconcile the observed duty cycle with smoothly-rising SFHs.  To this end, we have already demonstrated encouraging agreement between our predictions and the observed occupancy and $M_1/M_{\\mathrm{min}}$. A more detailed study of the CLF and its associated HOD in simulations with and without outflows would bring the impact of outflows on the HOD into sharper relief while providing a new class of template HODs for interpreting observations.  Such a study would involve using our predicted HOD to populate halos from a larger-volume N-body simulation in order to provide good statistics, a task that is well beyond the scope of our current work.  On the observational side, further progress will follow from a more systematic reconstruction of the high-redshift SFR-$\\mstar$ correlation and its evolution, which are the most important observational probes of the ``typical\" SFH in an epoch where the SFHs of individual galaxies remain unconstrained.  This involves work on three fronts: First, deeper rest-frame optical constraints on larger samples of galaxies are required in order to leverage existing UV-selected samples.  Second, progress requires improved SED modeling that invokes physically-motivated SFHs (rather than, for example, instantaneous bursts) and accounts for nebular emission in order to reduce systematic errors.  Third, uniform SED-fitting studies incorporating common modeling assumptions~\\citep[for example,][]{sch10} will reduce artificial scatter owing to differing modeling assumptions. Because of the importance of the SFR-$\\mstar$ plot, we heartily encourage the publication of inferred physical properties order to facilitate further theoretical comparisons.  In order to gain a better understanding of the scatter in the SFR-$\\mstar$ relation as well as the implications of the low observed halo occupancy, we encourage further inquiry into the overlap between UV-selected and optically-selected samples at $z\\geq4$, following the example of~\\citet{bra07}.  These studies are a strong test of our prediction that SFHs are predominantly smooth rather than bursty.  With the growing body of observational evidence that high-redshift SFHs are smooth rather than bursty and the implication that the progenitors of $V$-dropouts, for example, are visible as $i$-dropouts, it is important that SED-fitting studies account for galaxies' hypothetical progenitors at earlier epochs in addition to their observed fluxes.  For example, exponentially decaying SFHs with a decay time of 300 Myr applied to $V$-dropouts predict bright progenitor $i$-dropouts whose colours/SSFRs are observationally ruled out (unless they are very dusty, which again reinforces the need for study of optically-selected samples).  For the same reason, it is now becoming possible to constrain high-redshift galaxy SFHs using statistical methods applied at multiple epochs~\\citep[see, for example,][]{con09,mar10,pap10}.  Such studies provide complementary tests of our predicted SFHs.  Finally, we have outlined a number of predictions regarding the clustering properties of faint ($M_{1350}>-20$) high-redshift galaxies that may be amenable to study using current samples (for example, from GOODS) and that should be considered in the design of upcoming Hubble and Spitzer programs to study the distant universe."
    },
    "1005/1005.1196_arXiv.txt": {
        "abstract": "The diffuse meta-galactic radiation field at ultraviolet to infrared wavelengths - commonly labeled extragalactic background light (EBL) - contains the integrated emission history of the universe. Difficult to access via direct observations, indirect constraints on its density can be derived through observations of very-high energy (VHE; E$>$100\\,GeV) $\\gamma$-rays from distant sources: the VHE photons are attenuated via pair-production with the low energy photons from the EBL, leaving a distinct imprint in the VHE spectra measured on earth. Discoveries made with current generation VHE observatories like \\He and \\Ma enabled strong constraints on the density of the EBL, especially in the near-infrared. In this article the prospect of future VHE observatories to derive new constraints on the EBL density are discussed. To this end, results from current generation instruments will be extrapolated to the future experiment's sensitivity and investigated for their power to enable new methods and improved constraints on the EBL density. ",
        "introduction": "The observation of very-high energy $\\gamma$-rays (VHE; $E>100$\\,GeV) from distant sources offers the unique possibility to probe the density of the meta-galactic radiation field at ultraviolet (UV) to infrared (IR) wavelengths, which is commonly labeled the extragalactic background light (EBL; typically 0.1-100\\,$\\mu$m). The VHE $\\gamma$-rays interact with the low energy EBL photons via the pair production process ($\\gamma_{\\mathrm{VHE}} \\gamma_{\\mathrm{EBL}} \\rightarrow e^+ e^-$) and the flux is attenuated \\cite{nikishov:1962a,gould:1967a}. This attenuation can leave distinct signatures in the measured VHE spectra. With assumptions about the source physics and the spectrum emitted at the source location (intrinsic spectrum), constraints on the density of the EBL can be derived \\cite[e.g.][]{stecker:1994a, dwek:1994a}.  The current generation of VHE instruments (\\He, \\Ma, \\Ve) significantly increased the number of known extragalactic VHE sources from 4 in the year 2003 to more than 25 today.  These discoveries, combined with the advanced spectral resolution of these instruments and the wide energy range they cover, led to new strong constraints on the EBL density, in particular at optical to near-IR (NIR)  wavelengths \\cite{aharonian:2006:hess:ebl:nature,mazin:2007a,albert:2008:magic:3c279:science}. Since these limits depend on assumptions about the source physics, the strong constraints also sparked intense discussions on the validity of the assumptions and possible caveats \\cite[e.g.][]{aharonian:2002c, katarzynski:2006a, reimer:2007a,aharonian:2008a,krennrich:2008a,raue:2009b}. This discussion has not yet converged and there are interesting arguments for both sides.  Current generation systems have recently been upgraded (\\MaII) or the upgrades are under construction (\\HeII). These upgrades are mainly aimed to improve the overall sensitivity by a factor two to three and extend the energy range toward the lower energy regime of 20 to 100\\,GeV. This will lead to some improvements, but a quantitative difference or a breakthrough compared to the performance of the existing facilities will only be achieved with an order of magnitude improvement in sensitivity. The Next Generation Cherenkov Telescope Systems (NGCTS) are in the advanced planing phase aiming to achieve this order of magnitude improvement: the Cherenkov Telescope Array (CTA\\footnote{http://www.cta-observatory.org/}) \\cite{hermann:2007a} and the Advanced Gamma-ray Imaging System (AGIS\\footnote{http://www.agis-observatory.org/}) \\cite{buckley:2008a}. Whereas CTA envisions to improve the sensitivity over a wide energy range from the few tens of GeV to the multi TeV regime, AGIS mainly concentrates on energies above 100 GeV, an extended field of view and an improvement of the angular resolution.  The potential of these upcoming experiments to probe the EBL is the topic of this article. While an order of magnitude improvement in sensitivity for an astronomical instrument will always lead to new and unexpected results, this article will - as a first step - focus on known results and their extrapolation according to the sensitivities of the next generation instruments. Emphasis will be on new techniques enabled by the performance features (extended sensitivity and energy range) of the upcoming instruments.  For the calculations in the paper a standard $\\Lambda$CDM cosmology with $h=\\Omega_\\Lambda=0.7$ and $\\Omega_M=0.3$ is adopted. ",
        "conclusions": "In this paper the potential of a Next Generation Cherenkov Telescope System (NGCTS)  to study the EBL through observations of VHE spectra from distant sources is explored. In the focus of the study lies the energy range between 40\\,GeV and 10\\,TeV, where a factor 10 improvement in sensitivity over current generation experiments is expected. Two different methods are investigated:  (i) utilizing the unabsorbed part of the VHE spectrum in the energy range 40-100\\,GeV and (ii) searching for attenuation modulation signatures at energies between 100\\,GeV to 7\\,TeV. While some caveats, like e.g. the exact shape of the intrinsic spectra, do exist, overall the two methods show promising results, clearly go beyond what is possible with current generation instruments.  EBL attenuation is a global phenomenon that affects the spectra of all sources in the same way. It is therefore quite natural to expand the studies on the EBL by combining the results from different methods, sources or source flux states. This can e.g. be done by utilizing log-likelihood methods to combine the constraints from different sources and methods. In addition only the low and mid-VHEs have been investigated. At energies $>10$\\,TeV further signatures from EBL attenuation (\"cut-off\") are expected, which have been used extensively in the past to derive limits on the EBL density in the FIR.  In this paper only data from a single instrument - a NGCTS - and minimal assumptions about intrinsic spectrum have been used to derive constraints on the EBL density. AGNs are observed all across the electromagnetic spectrum and the theory predicts connections between the observations at different energies. A wealth of multi-wavelength information is therefore accessible, which - in combination with a theoretical model - can be used to constrain the spectrum at VHE. A first model-dependent approach to determine the EBL level using a consistent SED modeling of detected blazars was discussed recently by \\cite{mankuzhiyil:2010a}. While a detailed discussion of this topic is beyond the scope of this paper, as an example it should only be mentioned the power of combining the \\Fe observations at MeV to GeV energies with observations from an NGCTS: such observations will cover over 9 decades (!) in energy with high precision.  The sources under study show strong flux variability in the VHE band. The high sensitivity of a NGCTS will enable to study the energy spectrum on very short time-scales in great detail. The time resolution will enable stronger constraints on the theoretical modeling of the source spectra since the experimental data will be available on shorter scales than the relevant changes in the emission regions. This will further help to disentangle source intrinsic effects and EBL attenuation.  In this study only the VHE spectra of a few selected known sources have been investigated. The NGCTS will bring new discoveries. With the high sensitivity in the 20 to 100\\,GeV range it will be possible to detect fainter sources and sources at higher redshift. A (sufficiently) large sample of sources will serve a two-fold purpose: (1) help to improve the understanding of the intrinsic source physics and (2) enable statistical studies of the EBL effects, e.g. by investigating the EBL attenuation features versus redshift. This will also enable to not only probe the present day EBL but to study its evolution with redshift. How such a study can be performed with the help of an NGCTS will be the topic of a second paper."
    },
    "1005/1005.0791_arXiv.txt": {
        "abstract": "{} {We present an adaptation of the standard scenario of disk-galaxy formation to the concordant $\\Lambda$CDM cosmology aimed to derive analytical expressions for the scale length and rotation speed of present-day disks that form within four different, cosmologically motivated protogalactic dark matter halo-density profiles.} {We invoke a standard galaxy-formation model that includes virial equilibrium of spherical dark halos, specific angular momentum conservation during gas cooling, and adiabatic halo response to the gas inflow. The mean mass-fraction and mass-to-light ratio of the central stellar disk are treated as free parameters whose values are tuned to match the zero points of the observed size-luminosity and circular speed-luminosity relations of galaxies.} {We supply analytical formulas for the characteristic size and rotation speed of disks built inside Einasto $r^{1/6}$, Hernquist, Burkert, and Navarro-Frenk-White dark matter halos. These expressions match simultaneously the observed zero points and slopes of the different correlations that can be built in the $RV\\!L$ space of disk galaxies from plausible values of the galaxy- and star-formation efficiencies.} {} ",
        "introduction": "In the current hierarchical galaxy-formation paradigm disk-galaxies are born out of the hot gas-atmospheres associated with the potential well of virialized cold dark matter (CDM) halos. It is assumed that baryons have initially both the same density profile and specific angular momentum distribution as DM -- the latter achieved, for instance, through tidal interactions with neighboring objects in the precollapse phase \\citep[e.g.][]{Pee69}. As the gas radiates its energy it cools and starts to fall towards the center of the DM halo maintaining its specific angular momentum, where it settles into a rotationally supported disk. The assembly of a concentration of cold baryons at the bottom of the gravitational potential well on timescales longer than the free-fall time produces the adiabatic contraction of the dark halo\\footnote{In modern literature, the mode and amount of halo contraction are actually a matter of debate \\citep[e.g.][]{Dut07,Tis09}. The outcome, however, remains unchanged: the properties of disk galaxies are linked to those of their host halos.}. In this standard picture, the internal properties of disk galaxies are expected to be largely dictated by those of their host halos, and through the latter, by those of the background cosmology too.  Theoretical predictions for the distribution of disk galaxies in the space of disk scalelength (or size), fiducial (usually, maximum or asymptotic) rotational speed, and luminosity (or mass) based, partially or totally, on the scenario just outlined are abundant in the literature \\citetext{e.g., \\citealt*{MMW98}, hereafter \\citeauthor{MMW98}; \\citealt{Piz05,Dut07}}. They are widely used in semi-analytic cosmological models, pre-prepared numerical simulations of galaxy groups and clusters, and studies of disk-galaxy scaling relations.  While nowadays there are extensive and comprehensive investigations of the correlations between disk-galaxy properties that deal with the scatter and covariances of the variables and allow for different modes of halo contraction \\citep[e.g.][]{Dut07}, it is not always feasible to implement such sophisticated treatments whenever one needs to estimate the scaling of the basic structural and kinematic parameters of galaxies. The simplest alternative is the use of scaling laws derived directly from fits to a given set of observations. However, because of their lack of theoretical foundation, these formulas cannot be extrapolated to explain the properties of galaxies other than those from which they are derived. Halfway between these two options is the possibility of using analytical expressions endowed with a physical basis that enables their application to a wide range of galactic and halo parameters. It is precisely with this aim that we here introduce a self-consistent pure disk-formation model that follows the well-known approach by \\citeauthor{MMW98} adapted to the canonical $\\Lambda$CDM concordance cosmology and to four different mass-density distributions for the protogalactic dark halos. This updated scenario is capable of matching \\emph{simultaneously} with very good accuracy the zero points and slopes of the observed correlations in the $RV\\!L$ space of disk galaxies from reasonably realistic values of its input parameters. Yet its most valuable characteristic is its ease of implementation, as we approximated the model predictions for the scale length and rotation speed of disks by analytical expressions. The supplied equations can come in handy for situations that require the generation of large numbers of galaxies with intrinsic attributes in good agreement with the mean observed trends, especially when the relative abundances of these objects are known in advance. ",
        "conclusions": "We formulated a standard formation model of disk galaxies inside DM halos within the concordant $\\Lambda$CDM cosmology that simultaneously predicts with high accuracy the main trends of the observed fundamental scaling relations of nearby galaxies in $RV\\!L$ space. This modeling has been developed with the sole aim of deriving physically sound analytical expressions for predicting the central properties that characterize the light profiles and rotation curves of \\emph{typical spirals}. We supply formulas for Einasto $r^{1/6}$, Hernquist, Burkert, and Navarro-Frenk-White protogalactic halo mass density distributions that provide a similarly good overall description of the data on disks for realistic enough values of the model free parameters. We find that, for a given $\\ld$, the predictions of the Einasto $r^{1/6}$, Hernquist, and Navarro-Frenk-White profiles are relatively similar, while the Hernquist profile -- the only density law investigated that does not follow a $\\rho\\propto r^{-3}$ behavior near $\\Rvir$ -- requires values of $\\md$ and $\\MLd$ that are lower by about a factor of 0.70 and 0.85, respectively.    \\begin{figure*} \\centering \\begin{center} \\begin{tabular}{cc} \\resizebox{74mm}{!}{\\includegraphics{13458f01.eps}} & \\resizebox{74mm}{!}{\\includegraphics{13458f02.eps}} \\\\ \\resizebox{74mm}{!}{\\includegraphics{13458f03.eps}} & \\end{tabular}  \\caption{Scale relations for nearby disks. \\emph{Upper-left:} Disk-specific angular momentum as a function of $\\Vmax$. The cross shows the barycenter of the data. \\emph{Upper-right:} $I$-band TF relation. The dashed line shows the one derived by \\citet{Mas06} from SFI++ data \\citep{Spr07}. The cross shows the barycenter of the data. \\emph{Lower-left:} Central-disk surface-mass density vs.\\ $\\Vmax$. Squares with error bars show the median observational values and the first and third quartiles in each velocity bin. In all these plots the solid lines show model predictions for pure disks with $\\ld=0.04$ using the mean $c(\\Mvir)$ relation at $z=0$, while the dotted curves show from top to bottom the predictions resulting from adopting the 2.3th and 97.7th percentiles of the concentration distribution ($2\\sigma$ scatter) assuming $\\sigma_{\\log c}=0.11$ \\citep{Mac08}. The error boxes represent, for comparison, the values measured for the MW from: $\\Rd=2.5\\pm 0.5$ kpc, $\\Md=5\\pm 0.5\\times 10^{10}\\msun$, and $\\Vmax=220\\pm 20$\\kms\\ \\citep{BT08}, and $L_I=(4\\pm 0.8)\\times 10^{10}\\lsun$ \\citep{Por07}. The data clouds are build on the compilation of $I$-band absolute magnitudes and \\hi\\ rotation widths by \\citet{Spr07} and on the $I$-band disk scalelength measurements by \\citet{Cou07}. \\label{figallSB}} \\end{center} \\end{figure*}    The reader may have noticed that our best models yield for $\\MLd$, i.e. for the inverse of the average star-formation efficiency, values somewhat lower than those inferred from population synthesis calculations \\citep[e.g.][]{Piz05}. We stress however that the observational estimates of this parameter are affected by considerable uncertainties, our prediction that the average mass-to-light ratio of disks is $\\sim 1 (M/L_I)_\\odot$, which is consistent with submaximal disks arguments \\citep{CR99,KdN08}, as well as relatively close to the values adopted as input in more sophisticated models of disk formation \\citep{Dut07}. On the other hand, we find that the predicted values of $\\md$ are directly correlated with those adopted for $\\ld$. In particular we note that a value of $\\ld=0.03$, which coincides with the median of the distribution of the spin parameter for relaxed halos derived by \\citet{Mac08}, implies a small current average galaxy-formation efficiency, $\\md/\\fb<0.2$. This agrees well with the predictions of galaxy evolution from halo occupation models \\citep{ZCZ07} and methods that match the stellar mass function to that of the halo \\citep{CW09}. Further recent support for low $\\md$ (and $\\ld$, according to our model) comes for instance from weak lensing measurements \\citep{Man06} and from the roughly universal distributions of this parameter obtained by \\citet{Sal09} for various implementations of feedback in large cosmological $N$-body/gasdynamical simulations. Notice also the fifth column in Table~\\ref{parameters}, where we list the ratio $\\Vmax/\\Vvir$ calculated for a MW-mass halo, which increases with increasing $\\ld$ and decreasing $\\Mvir$. As stated by \\citet{Dut07}, the relatively high values we predict for this ratio -- a characteristic common to standard models -- would likely hamper a simultaneous match to the galaxy LF that, according to semi-analytical models of galaxy formation, requires the condition $\\Vmax\\sim\\Vvir$.  We made no attempt to explore the scatter of the observed scaling relations and the covariance that exists between model parameters, except for Fig.~\\ref{figallSB}, where we carry out a naive comparison between the spread of the data and that resulting from taking into account the predicted scale of the probability distribution of the halo concentration. Including this and other sources of scatter, such as the variance of the halo spin parameter, or the dependence of the concentration-mass relation on the adopted cosmology \\citep[e.g.][]{Mac08}, would undoubtedly enrich the analysis. Yet, a thorough investigation of scatter requires dealing with the joint probability distribution of all the parameters entering the model and, in particular, with all their covariances (not just the variances), which ideally should be corrected from measurement errors. This far exceeds the scope of our present research. We note in addition that efforts in the direction just outlined will soon be much more effective when they can be applied to objective, homogeneous, and complete $RV\\!L$ datasets free of nontrivial selection biases, as those build from the cross-correlation of wide-area spectroscopic optical and \\hi\\ surveys (e.g., Toribio et al.\\ 2010, in preparation).  Finally, we wish to comment on the possibility of extending our model predictions to distant galaxies by adopting a $c(\\Mvir)$ relationship of the form \\begin{equation}\\label{c-M(z)} c(\\Mvir,z)=c(\\Mvir,0)g(z)^{-1}\\;, \\end{equation} with $c(\\Mvir,0)$ given in Eq.~(\\ref{c-M}) and $g(z)$ the concentration growth factor, which can be taken proportional to $H(z)^{2/3}$, as found in a recent modification of the original \\citet{Bul01} model for \\emph{WMAP} cosmologies by \\citet{Mac08}."
    },
    "1005/1005.2331_arXiv.txt": {
        "abstract": "{We present {\\it Herschel}-SPIRE observations of the perturbed galaxy NGC4438 in the Virgo cluster. These images reveal the presence of extra-planar dust up to $\\sim$4-5 kpc away from the galaxy's disk. The dust closely follows the distribution of the stripped atomic and molecular hydrogen, supporting the idea that gas and dust are perturbed in a similar fashion by the cluster environment. Interestingly, the extra-planar dust lacks a warm temperature component when compared to the material still present in the disk, explaining why it was missed by previous far-infrared investigations. Our study provides evidence for dust stripping in clusters of galaxies and illustrates the potential of {\\it Herschel} data for our understanding of environmental effects on galaxy evolution. } ",
        "introduction": "Clusters of galaxies are extremely hostile environments for star-forming galaxies. A plethora of observations and numerical simulations have revealed how gravitational and hydrodynamical interactions can affect cluster spirals (e.g., \\citealp{review}): stars and gas can be stripped, galaxy morphologies can be changed and star formation activity can be enhanced and/or quenched. However, still very little is known about the effects of the environment on the dust content of cluster galaxies. Since the dust is mixed with the interstellar medium (ISM), the common expectation is that the environment should affect the dust in a similar fashion as the gas. However this hypothesis has still to be confirmed observationally. The launch of {\\it Herschel} \\citep{pilbratt10} has opened a new era in the study of dust in galaxies. Thanks to its high spatial resolution and sensitivity to all dust components, {\\it Herschel} should be able to determine whether dust is stripped from infalling cluster spirals and dispersed into the intra-cluster medium (ICM).  One of the most dramatic examples of the effects of the environment on nearby galaxies is represented by the disturbed galaxy NGC4438, in the Virgo cluster. NGC 4438 is a highly H{\\sc i}-deficient early-type spiral showing stellar tails \\citep{kenney95,n4438} and extra-planar gas \\citep{vollmer09}. Numerical simulations show that only a combination of ram-pressure stripping and tidal interactions is able to reproduce the disturbed morphology and kinematics of NGC4438 \\citep{vollmer05}. However, while \\cite{combes88} and \\cite{vollmer05} invoked a high-velocity ($\\sim$800 \\kms) gravitational interaction with the companion early-type galaxy NGC4435, \\cite{kenney08} have revealed the presence of a series of H$\\alpha$+[NII] filaments connecting NGC4438 to the giant elliptical M86 (see Fig.~\\ref{spiremap}). These new observations favour a more complex scenario in which NGC4438 has recently ($\\sim$100 Myr ago) interacted with M86 \\citep{kenney08}. Although the detailed history of NGC4438 is still unclear, its peculiar properties make it an ideal target to test the power of {\\it Herschel} in unveiling the effects of the environment on the dust properties of cluster galaxies.  In this Letter we present {\\it Herschel}-SPIRE \\citep{griffin10} observations of NGC4438 obtained as part of the {\\it Herschel} Science Demonstration (SD) Phase. The observations of the elliptical galaxy M86 are presented in a companion paper \\citep{gomez10}. We assume for NGC4438 a distance of 17 Mpc \\citep{gav99}, corresponding to a linear scale of 82 pc/\\arcsec.  \\begin{figure*} \\centering \\includegraphics[width=15.5cm]{14547fg1.ps} \\caption{The NGC4438/M86 region as seen by SPIRE. This mosaic is obtained by combining the 250$\\mu$m maps centered on NGC4438 and M86 \\citep{gomez10}. The blue contours show the extended H$\\alpha$+[NII] filaments discovered by \\cite{kenney08}. Note that the intensity excess south of NGC4438 (near the edge of the frame) is just a residual of the baseline subtraction.} \\label{spiremap} \\end{figure*} ",
        "conclusions": "The SPIRE data alone are not sufficient to determine whether the extra-planar dust is in the process of being removed from NGC4438 or, for example, falling back onto the disk. However, by combining multiwavelength observations with detailed numerical simulations, \\cite{vollmer05,vollmer09} have shown that the distribution and kinematics of the different components of the ISM in NGC4438 can only be reproduced via a combination of tidal interaction and ram-pressure stripping. Although the details of the gravitational interaction are still unclear (i.e., M86 and/or NGC4435, \\citealp{vollmer09b}), it appears that strong on-going ram-pressure (in addition to tidal forces) is necessary to reproduce the properties of the extra-planar gas component in knot B. In this case, whatever the exact mechanisms affecting NGC4438, it is likely that at least part of the dust in the west side of the galaxy is in the process of being removed from the disk thus providing evidence of dust stripping by environmental effects. If completely removed, the stripped dust will be dispersed in the ICM contributing to its metal enrichment.   Since in galaxies the dust is associated with the gaseous component of the ISM, it is generally expected that when the gas is stripped part of the dust can be removed as well. What is still unclear is how much dust follows the fate of the stripped hydrogen. Unfortunately, the SPIRE fluxes alone are not sufficient to accurately quantify the amount and temperature of the dust in knot B. Nevertheless, we can at least try to estimate the dust mass by using the fluxes at 250 and 350 $\\mu$m, where knot B is clearly resolved, and assuming $M_{dust} = (f_{\\nu} D^{2})/[\\kappa_{\\nu} B(\\nu, T)]$, where $f_{\\nu}$ is the flux density ($\\sim$0.98 and $\\sim$0.46 Jy at 250 and 350 $\\mu$m, respectively), $D$ is the distance and $\\kappa_{\\nu}$ is the absorption cross section per mass of dust. We adopt $\\kappa_{\\nu}$=4 and 1.9 cm$^2$ g$^{-1}$  at 250 and 350 $\\mu$m respectively \\citep{draine03ara}. The total mass of dust in knot B is in the range $\\sim$2$\\times$10$^{6}$-2$\\times$10$^{7}$ M$_{\\odot}$ for a dust temperature between 10 and 20 K. In order to determine the total gas mass in the extra-planar cloud, we combined the estimates of M(H$_{2}$)$\\sim$4.7$\\times$10$^{8}$ M$_{\\odot}$ (including helium) and M(H{\\sc i}) $\\sim$1.5$\\times$10$^{8}$ M$_{\\odot}$ obtained by \\cite{vollmer05} and \\cite{hota07}, respectively. The total gas-to-dust ratio of knot B is in the range $\\sim$30-300, not significantly different from the values observed in nearby galaxies (e.g., \\citealp{draine07,galliano08}) and across the disk of star-forming spirals (e.g., \\citealp{mateos09,bendo10b,pohlen10}). This seems to support a scenario in which a cloud of gas does not lose a significant amount of its heavy elements when removed from the disk. It is in fact quite remarkable how well the cold dust follows the spatial distribution of the cold hydrogen across the whole galaxy, even in such a highly perturbed system like NGC4438.  Interestingly, the high degree of spatial correlation between cold dust and cold hydrogen in NGC4438 is different from what is observed by {\\it Herschel} in M86 \\citep{gomez10}. In M86 the dust is mainly associated with the H$\\alpha$+[NII] emitting gas and only mildly correlated with the cold atomic component. This is intriguing and might suggest different dust properties (e.g., stripping mechanism, heating source, etc.) in the two galaxies.  Finally, we note that it is unclear whether the extra-planar dust is missing a hot component just because the interstellar radiation field outside the plane is not strong enough to keep it hot, or if the cloud was already dominated by a cold dust component while in the disk. Although the $f(250)/f(350)$ flux ratio is lower than the value observed in the central part of NGC4438, it is not significantly different from what is observed in the outer parts of the disk of M81 \\citep{bendo10}, M99 and M100 \\citep{pohlen10}.  In summary, in this paper we provide evidence for dust stripping by environmental effects in NGC4438. The high spatial resolution and sensitivity to cold dust of the SPIRE camera allowed us to discover an extra-planar cold dust component completely missed by previous far-infrared surveys. The strong spatial correlation between cold dust, atomic and molecular hydrogen suggests that when the gas is removed also the dust is pulled out the galactic disk. These results provide interesting insights into the evolution of an extreme case among perturbed cluster galaxies. Once combined with the discovery of truncated dust disks in H{\\sc i}-deficient Virgo cluster spirals \\citep{cortese10b}, our analysis clearly highlights the great potential of {\\it Herschel} for our understanding of the effects of the cluster environment on the dust properties of galaxies."
    },
    "1005/1005.3101_arXiv.txt": {
        "abstract": "Braneworld models with variable brane tension $\\lambda $ introduce a new degree of freedom that allows for evolving gravitational and cosmological constants, the latter being a natural candidate for dark energy.  We consider a thermodynamic interpretation of the varying brane tension models,  by showing that the field equations with variable $\\lambda $ can be interpreted as describing matter creation in a cosmological framework. The particle creation rate is determined by the variation rate of the brane tension, as well as by the brane-bulk energy-matter transfer rate. We investigate the effect of a variable brane tension on the cosmological evolution of the Universe, in the framework of a particular model in which the brane tension is an exponentially dependent function of the scale factor. The resulting cosmology shows the presence of an initial inflationary expansion, followed by a decelerating phase, and by a smooth transition towards a late accelerated de Sitter type expansion. The varying brane tension is also responsible for the generation of the matter in the Universe (reheating period). The physical constraints on the model parameters, resulted from the observational cosmological data, are also investigated. ",
        "introduction": "The idea of embedding our Universe in a higher dimensional space has attracted a considerable interest recently, due to the proposal by Randall and Sundrum that our four-dimensional (4D) spacetime is a three-brane, embedded in a 5D spacetime (the bulk) ~\\cite{Randall,Randall2}. This proposal is based on early studies on superstring theory and M-theory, which have suggested that our four dimensional world is embedded into a higher dimensional spacetime. Particularly, the 10 dimensional $E_8 \\otimes E_8$ heterotic superstring theory is a low-energy limit of the 11 dimensional supergravity, under the compactification scheme $M^{10}\\times S_1 / Z_2$ \\cite{Witten,Witten2}. Thus, the 10 dimensional spacetime is compactified as $M^4 \\times CY^6 \\times S_1 / Z_2$, implying that our Universe (a brane) is embedded into a higher dimensional bulk. In this paradigm, the standard model particles are open strings, confined on the braneworld, whilst the gravitons and the closed strings can freely propagate into the bulk \\cite {Polchinski}.  The Randall-Sundrum Type II model has the virtue of providing a new type of compactification of gravity \\cite{Randall,Randall2}. Standard 4D gravity can be recovered in the low-energy limit of the model, with a 3-brane of positive tension embedded in 5D anti-de Sitter bulk. The covariant formulation of the braneworld models has been formulated in \\cite{Shiromizu}% , leading to the modification of the standard Friedmann equations on the brane. It turns out that the dynamics of the early Universe is altered by the quadratic terms in the energy density and by the contribution of the components of the bulk Weyl tensor, which both give a contribution in the energy momentum tensor. This implies a modification of the basic equations describing the cosmological and astrophysical dynamics, which has been extensively considered recently \\cite{all2}.  The recent observations of the CMB anisotropy by WMAP \\cite{Spergel} have provided convincing evidence for the inflationary paradigm \\cite{Guth}, according to which in its very early stages the Universe experienced an accelerated (de Sitter) expansionary phase (for recent reviews on inflation see \\cite{infl}).  At the end of inflation, the Universe is in a cold and low-entropy phase, which is utterly different from the present hot high-entropy Universe. Therefore the Universe should be reheated, or defrosted, to a high enough temperature, in order to recover the standard Hot Big Bang \\cite{reh}. The reheating process may be envisioned as follows: the energy density in zero-momentum mode of the scalar field decays into normal particles with decay rate $\\Gamma$. The decay products then scattered and thermalize to form a plasma \\cite{infl}.  Apart from the behavior of the inflaton field, the evolutions of dark energy and dark matter in reheating stage were also considered. In \\cite{Susperregi}% , dark energy and dark matter were originated from a scalar field in different stages of the inflation, according to a special form of potential. Meanwhile, the conditions for unifying the description of inflation, dark matter and dark energy were considered in \\cite{Liddle3}. A specific model was later proposed in \\cite{Cardenas}, by using a modified quadratic scalar potential. The candidates of dark matter in \\cite{Liddle3} and \\cite {Cardenas} were oscillations of a scalar field. However, it may be possible that dark matter existed on its own without originating from the scalar field. This may pose less stringent constraint on the scalar field, so that dark matter can be included in inflation paradigm in a easier way. On the other hand, it was proposed that the decay products of scalar field acquired thermal mass \\cite{Kolb}.  The reheating in the braneworld models has also been considered recently. In the context of the braneworld inflation driven by a bulk scalar field, the energy dissipation from the bulk scalar field into the matter on the brane was studied in \\cite{HiTa03}. The obtained results supports the idea that the brane inflation model, caused by a bulk scalar field, may be a viable alternative scenario of the early Universe. The inflation and reheating in a braneworld model derived from Type IIA string theory was studied in \\cite {BrDa03}. In this model the inflaton can decay into scalar and spinor particles, thus reheating the Universe. A model in which high energy brane corrections allow a single scalar field to describe inflation at early epochs and quintessence at late times was discussed in \\cite{SaDaSh03}. The reheating mechanism in the model originates from Born-Infeld matter, whose energy density mimics cosmological constant at very early times and manifests itself as radiation subsequently. The particle production at the collision of two domain walls in a 5-dimensional Minkowski spacetime was studied in \\cite{TaMa04}. This may provide the reheating mechanism of an ekpyrotic (or cyclic) brane Universe, in which two BPS branes collide and evolve into a hot big bang Universe. The reheating temperature $T_{\\rm RH}$ in models in which the Universe exits reheating at temperatures in the MeV regime was studied in \\cite{Hannestad}, and a minimum bound on $T_{\\rm RH}$ was obtained. The derived lower bound on the reheating temperature also leads to very stringent bounds on the compactification scale in models with $n$ large extra dimensions. The dark matter problem in the Randall-Sundrum type II braneworld scenario was discussed in \\cite{Pa05}, by assuming that the lightest supersymmetric particle is the axino. The axinos can play the role of cold dark matter in the Universe, due to the higher reheating temperatures in the braneworld model, as compared to the conventional four-dimensional cosmology. The impact of the non-conventional brane cosmology on the relic abundance of non-relativistic stable particles in high and low reheating scenarios was investigated in \\cite{DaKh06}. In the case of high reheating temperatures, the brane cosmology may enhance the dark matter relic density by many order of magnitudes, and a stringent lower bound on the five dimensional scale may be obtained. In the non-equilibrium case, the resulting relic density is very small. The curvaton dynamics in brane-world cosmologies was studied in \\cite{PaZa06}.   Brane-worlds with non-constant tension, based on the analogy with fluid membranes, which exhibit a temperature-dependence according to the empirical law established by E\\\"otv\\\"os, were introduced in \\cite{Ger1}. This new degree of freedom allows for evolving gravitational and cosmological constants, the latter being a natural candidate for dark energy. The covariant dynamics on a brane with variable tension was studied in its full generality, by considering asymmetrically embedded branes, and allowing for non-standard model fields in the 5-dimensional space-time. This formalism was applied for a perfect fluid on a Friedmann brane, which is embedded in a 5-dimensional charged Vaidya-Anti de Sitter space-time. For cosmological branes a variable brane tension leads to several important consequences. A variable brane tension may remove the initial singularity of the Universe, since the brane Universe was created at a finite temperature $T_{c}$ and scale factor $a_{\\min}$ \\cite{Ger2}. Both the brane tension and the 4-dimensional gravitational coupling 'constant' increase with the scale factor from zero to asymptotic values. The 4-dimensional cosmological constant is dynamical, evolving with $a$, starting with a huge negative value, passing through zero, and finally reaching a small positive value. Such a scale--factor dependent cosmological constant has the potential to generate additional attraction at small $a$ (as dark matter does) and late-time repulsion at large $a$ (dark energy). The evolution of the brane tension is compensated by energy interchange between the brane and the fifth dimension, such that the continuity equation holds for the cosmological fluid \\cite{Ger2}. The resulting cosmology closely mimics the standard model at late times, a decelerated phase being followed by an accelerated expansion. The energy absorption of the brane drives the 5D space-time towards maximal symmetry, thus becoming Anti de Sitter. Other physical and cosmological implications of a varying brane tension have been considered in \\cite{yun}.    It is the purpose of the present paper to further investigate the cosmological implications of a varying brane tension. As a first step in our study, we consider a thermodynamic interpretation of the varying brane tension models,  by showing that the field equations with variable $\\lambda $ can be interpreted as describing matter creation in a cosmological framework. The particle creation rate is determined by the variation rate of the brane tension, as well as by the brane-bulk energy-matter transfer rate. In particular, by adopting a theoretical model in which the brane tension is a simple function of the scale factor of the Universe, we consider the possibility that the early inflationary era in the evolution of the brane Universe was driven by a varying brane tension. A varying brane tension may also be responsible for the generation of the matter after reheating, as well as for the late time acceleration of the Universe.  The present paper is organized as follows. In Section~\\ref{geo} we present the field equations of the brane world models with varying brane tension and we write down the basic equations describing the cosmological dynamics of a flat Friedmann-Robertson-Walker Universe. The thermodynamic interpretation of the brane-world models with varying brane tension and brane-bulk matter-energy exchange is considered in Section~\\ref{therm}. A power-law inflationary brane-world model with varying brane tension and non-zero bulk pressure is obtained in Section~\\ref{pow}. The analytical behavior of the cosmological model with varying brane tension is considered in Section~\\ref{scale}, by using the small and large time approximations for the brane tension. The numerical analysis of the model is performed in Section~\\ref{numerical}. We discuss and conclude our results in Section~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  In the present paper we have considered some cosmological implications of the braneworld models with variable tension. We have considered a thermodynamic interpretation of the model, and we have shown that from a thermodynamic point of view a variable brane tension can describe particle production processes in the early Universe, as well as the entropy production in the early stages of the cosmological evolution. A simple power-law inflationary cosmological model has also been obtained.   By adopting a simple analytical form for $\\lambda$, we have obtained a complete description of the dynamics and evolution of the Universe from the stage of the inflation to the phase of the late acceleration. Moreover, the differences between the energy scales of the theoretical vacuum energy at inflationary epoch and during late time acceleration are studied.  A variable brane tension can drive the inflationary evolution of the Universe, and also be responsible for matter creation in the post-inflationary phase (the reheating of the Universe). In the model we have adopted the brane tension converges to a constant that gives the gravitational constant, and the residue after cancelation with the $\\Lambda_5$ term would give the dark energy. It is interesting to note that, as opposed to the standard cosmological scenarios, in the present model matter creation takes place during the entire inflationary phase, but reaches a maximum after the exponential expansion of the Universe ends. Therefore in the braneworld models with varying brane tension there can be no clear distinction between inflation and reheating. By using numerical analysis as well as approximate analytical methods, we have obtained the result that in this variable tension model there are 5 phases in the cosmological evolution of the Universe. At the beginning, there are the inflationary and the reheating phases. During these phases, the brane tension is the dominant energy in the Universe, and its magnitude is of the order of the vacuum energy at GUT scale. As the brane tension decays, the matter is created, and the Universe enters in the hot Universe phase of the Big Bang picture. The third phase is the quadratic density domination phase. This is a unique characteristic of the brane world scenario, during which the Universe is dominated by a quadratic term in the energy density of the radiation. The deceleration parameter increases from $-1$ during inflation to $5$ at this phase. To study the cosmological dynamics we have introduced a set of dimensionless quantities, which describe the evolution of the scaled densities in terms of a dimensionless time parameter $\\tau$. By using the results of the numerical simulations for the rescaled variables, one can obtain some constraints on the physical parameters of the model."
    },
    "1005/1005.4517_arXiv.txt": {
        "abstract": "{The far-IR/sub-mm spectral mapping facility provided by the \\textit{Herschel}-PACS and HIFI instruments has made it possible to obtain, for the first time, images of \\wat\\ emission with a spatial resolution comparable to ground based mm/sub-mm observations.} {In the framework of the Water in Star-forming regions with \\textit{Herschel} (WISH) key program, maps in water lines of several outflows from young stars are being obtained, to study the water production in shocks and its role in the outflow cooling. This paper reports the first results of this program, presenting a PACS map of the o-\\wat\\ 179~\\um\\, transition obtained toward the young outflow L1157.} {The 179~\\um\\, map is compared with those of other important shock tracers, and with previous single-pointing ISO, SWAS, and Odin water observations of the same source that allow us to constrain the \\wat\\ abundance and total cooling. } {Strong \\wat\\, peaks are localized on both shocked emission knots and the central source position. The \\wat\\, 179~\\um\\, emission is spatially correlated with emission from \\h\\ rotational lines, excited in shocks leading to a significant enhancement of the water abundance. Water emission peaks along the outflow also correlate with peaks of other shock-produced molecular species, such as SiO and NH$_3$. A strong \\wat\\,peak is also observed at the location of the proto-star, where none of the other molecules have significant emission. The absolute 179~\\um\\, intensity and its intensity ratio to the \\wat\\ 557~GHz line previously observed with Odin/SWAS indicate that the water emission originates in warm compact clumps, spatially unresolved by PACS, having a \\wat\\, abundance of the order of 10$^{-4}$. This testifies that the clumps have been heated for a time long enough to allow the conversion of almost all the available gas-phase oxygen into water. The total \\wat\\ cooling is $\\sim$\\,10$^{-1}$\\, L$_\\odot$, about 40\\% of the cooling due to \\h\\ and 23\\% of the total energy released in shocks along the L1157 outflow. } {} ",
        "introduction": "Among the main coolants in molecular shocks, water is the tracer most sensitive to physical variations and the temporal evolution of protostellar outflows, thus representing a very powerful probe of their shock conditions and thermal history (e.g., Bergin et al. 1998). Water emission and excitation in shocks were studied extensively for the first time with the \\textit{Infrared Space Observatory} (ISO), the first space facility with spectroscopic capabilities in the mid- and far-IR. ISO surveyed the water emission in a large sample of outflows from young stellar objects (YSOs), providing  a global statistical picture of the importance of water in the outflow cooling and of variations in its abundance with shock properties and ages (see e.g., Nisini 2003, van Dishoeck 2004). Following ISO, the SWAS and Odin facilities made it possible to observe the ortho-\\wat fundamental line at 557~GHz, providing important constraints on the water abundance and kinematics in the cold outflow  gas components (e.g., Franklin et al. 2008, Bjerkeli et al. 2009). All these facilities, however, had poor spatial resolution (i.e. greater than 80$\\arcsec$), which did not allow one to locate the origin of the water emission nor study variations in abundances and excitation within individual flows.  In this framework, a sample of YSO outflows will be surveyed in different water lines by the PACS and HIFI instruments onboard the \\textit{Herschel} satellite, as part of the key program WISH (Water In Star-forming-regions with \\textit{Herschel}\\footnote{http://www.strw.leidenuniv.nl/WISH/.}). This paper presents the first results obtained from this survey, consisting of a PACS map of the H$_{2}$O 2$_{12}$-1$_{01}$ 179~\\um\\ line covering the outflow of the protostar L1157-mm, obtained during the \\textit{Herschel} Science Demonstration Phase. The 179~\\um\\ line is the transition connecting the lower two back-bone levels of ortho-\\wat. It is therefore one of the brightest water lines expected in collisionally excited conditions, thus representing an ideal tracer of the water distribution in shocked regions. L1157 is a well known outflow driven by a low mass class 0 object  (L1157-mm, $ L_{bol} \\sim 8.3$ L$_\\odot$, D=440 pc, Froebrich 2005). It is considered to be the prototype of chemically active flows, given the large number of different species detected in its shocked regions (e.g., Bachiller \\& Perez-Gutierrez 1997). This paper is presenting the first of several observations planned for this source by the WISH team. ",
        "conclusions": "We have presented a PACS spectral map of the \\wat\\ 179~\\um\\ transition obtained toward the L1157 protostellar outflow. Strong water emission peaks have been found at the location of previously-known shocked spots and correlate well with \\h\\ mid-IR rotational lines, as well as other important shock tracers, such as SiO and NH$_3$. The absolute 179~\\um\\, intensity and the intensity ratios with respect to the previously-observed 557~GHz line, indicate that the water emission originates in warm compact clumps, spatially unresolved by PACS, that have a \\wat\\, abundance of the order of 10$^{-4}$. The total \\wat\\ cooling has been estimated to be of the order of 8-9\\,10$^{-2}$L$_\\odot$, representing about 40\\% of the cooling due to \\h\\ and 23\\% of the total energy released in shocks along the L1157 outflow.  Additional \\textit{Herschel} PACS/HIFI observations of the L1157 outflow are planned  by the WISH program. These will enable us to investigate variations in the  water abundance within the outflow and correlate these with kinematical information."
    },
    "1005/1005.1618_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro}   \\subsection{Gyroscopic pumping} \\label{sec:introgyro}  Recently, Garaud \\& Acevedo-Arreguin (2009) (GAA09 hereafter) presented a preliminary analysis of a ``new'' mechanism for rotational mixing in stellar interiors and conjectured on its potential role in the depletion of Lithium in young Main Sequence stars. This mechanism, called ``gyroscopic pumping'', was originally studied in the context of Earth's atmospheric dynamics by Haynes et al. (1991) and later discussed in the astrophysical context by Gough \\& McIntyre (1998) and McIntyre (2007) who argued that it plays an important role in the dynamics of the solar interior. Loosely speaking, gyroscopic pumping occurs in any rotating fluid in the presence of additional stresses or forces which perpetually accelerate or decelerate the flows in the azimuthal direction. By conservation of angular momentum, the accelerated parcels of fluid move away from the rotation axis while the decelerated ones move toward the rotation axis, thus generating meridional fluid motion.  Stars with outer convective regions often exhibit a significant amount of surface differential rotation (e.g. Collier-Cameron 2007; Reiners 2007). This differential rotation is thought to be maintained by anisotropic and spatially varying Reynolds stresses (see R\\\"udiger 1989, for example), which tend to continually accelerate the equatorial regions and decelerate the poles in a manner most remarkably observed in the solar convection zone (Schou et al. 1998). If one assumes that the star is close to dynamical equilibrium, its mean rotation rate $\\Omega_{\\star}$ lies in between the polar and equatorial rotation rates. When viewed in a frame rotating with angular velocity $\\Omega_{\\star}$, the differential rotation of the star's convective zone forms a large-scale azimuthal flow pattern, typically prograde in the equatorial region and retrograde in the polar regions. The aforementioned ``gyroscopic pumping'' can then be viewed in two equivalent ways. As described earlier the constant acceleration of the equatorial regions is a local source of angular-momentum to the fluid, which by angular-momentum conservation must move outward from the rotation axis. Similarly, fluid in the polar regions must move toward the rotation axis. Alternatively, one may simply note that the Coriolis force associated with the azimuthal flows described above pushes the fluid away from the rotation axis in prograde regions, and toward the rotation axis in retrograde regions. As illustrated in Figure \\ref{fig:gyro}, the process naturally drives large-scale meridional flows throughout the outer convection region, with an upwelling near the equator, a poleward velocity near the surface and downwelling near the poles. The polar downwelling does not need to stop at the radiative--convective interface, and could in principle cause significant non-local mixing between the outer convection zone and the regions below.  A similar gyroscopic pumping process is likely to drive fluid motion from within an inner convective zone as well. While the internal rotation profile of stars with convective cores has never been observed, one can readily expect some degree of differential rotation since the turbulent stresses associated with the rotationally constrained convective motions are likely to act in a similar fashion to those of the outer convective region. In that case again, there is no a-priori reason for the flows thus generated to stop at the interface with the overlying radiative zone, and one may wonder how much mixing they induce in the star.  \\begin{figure}[h] \\centerline{\\epsfig{file=f1.epsf,width=8cm}} \\caption{Illustration of gyroscopic pumping in a star with an outer convection zone (shaded area). Slower rotating regions near the poles pump the fluid towards the rotation axis through the action of the Coriolis force (${\\bf F}_{\\rm c}$) while the rapidly rotating region near the equator pump the fluid away from the rotation axis. A large-scale circulation is thus driven, poleward near the surface. The position of the return flow depends on the radiative zone physics (see GAA09 and \\S\\ref{sec:intromix} for detail). } \\label{fig:gyro} \\end{figure}  \\subsection{Mixing in radiative regions induced by gyroscopic pumping} \\label{sec:intromix}  The first quantitative study of gyroscopic pumping in the context of stellar interiors was recently presented by GAA09. They focused on the solar case, i.e. a star with an outer convective region and an inner radiative region, and worked in the Boussinesq approximation (arguing that the solar radiative zone does not span too many pressure and density scaleheights). They showed that the fate of gyroscopically pumped flows -- how much overall mixing they induce beyond the convective zone -- depends equally on the thermal stratification and on the dynamical properties of the nearby radiative region.  In accordance with the earlier results of Garaud \\& Brummell (2008), they found that in stratified, rotating stars in quasi-steady dynamical balance, the effect of stratification on large-scale meridional flows is principally controlled by the quantity \\begin{equation} \\sigma = \\sqrt{\\rm Pr} \\frac{N}{\\Omega_\\star} \\mbox{  , } \\label{eq:sigmadef} \\end{equation} where Pr $=\\nu/\\kappa$ is the Prandtl number (where $\\nu$ is the local viscosity, and $\\kappa$ is the local thermal diffusivity), $N$ is the local Brunt-V\\\"ais\\\"al\\\"a -- or buoyancy -- frequency and $\\Omega_\\star$ is the mean stellar rotation rate. When $\\sigma$ is large, the effect of thermal stratification is strongly felt by the gyroscopically pumped meridional flows, which are exponentially damped away from the radiative--convective interface on the lengthscale $R_\\star / \\sigma$. Correspondingly, if  $\\sigma$ is small the flows can {\\it in principle} penetrate much more deeply into the radiative interior, and cause significant large-scale mixing.  It is important to note that $\\sigma$ depends on $\\Omega_\\star$, so that the effect of stratification (in the sense defined above) is strongly reduced for rapidly rotating stars. Figure \\ref{fig:sigma} shows an estimate of $\\sigma$ for various stars in the mass range $1.3M_\\odot-1.5M_\\odot$ at age 300Myr. In all cases, $\\sigma$ remains well-below unity showing that large-scale mixing of the radiative zone by gyroscopic pumping could be significant for these stars. \\begin{figure}[h] \\centerline{\\epsfig{file=f2.eps,width=8cm}} \\caption{Profile of $\\sigma(r)$ for three different stellar masses at age 300Myr. The 1.3$M_{\\odot}$, 1.4$M_{\\odot}$ and 1.5$M_{\\odot}$ stars are assumed to be rotating with the velocities $8 \\times 10^{-5}$rad/s, $1.3 \\times 10^{-4}$rad/s, and $1.5 \\times 10^{-4}$rad/s respectively. The stellar models are computed using a standard stellar evolution code developed by Bodenheimer et al. (2007).} \\label{fig:sigma} \\end{figure}  Crucially, however, GAA09 showed that even in the very weakly stratified limit ($\\sigma \\ll 1$), a second condition needs to be satisfied for the pumped flows to penetrate into the nearby radiative region. Indeed, this limit corresponds to the case where the system's dynamics are dominated by the balance between the perturbation to the pressure gradient and the Coriolis force. This well-known situation is called the ``Taylor-Proudman'' state, and (in the Boussinesq approximation) implies that all components of the velocity field are strongly constrained to be constant along the rotation axis. GAA09 showed that the Taylor-Proudman constraint can prohibit flow generated within the outer convection zone from entering the underlying radiative zone. Indeed, by mass conservation, any flows entering the radiative zone must somehow return to the convection zone. But such return flow would necessarily require breaking away from the Taylor-Proudman state. Hence, if there exists a region within the radiative zone where the Taylor-Proudman constraint is broken, then this region provides a channel through which the pumped flows can return, as illustrated in Figure \\ref{fig:gyro2}. If such a region does not exist, the meridional flows instead return within the convection zone causing negligible mixing in the radiative zone.  Since the Taylor-Proudman constraint is broken whenever there exist additional stresses of amplitude comparable with the Coriolis force, various mechanisms can be invoked. Of particular interest are Maxwell stresses in the presence of small- or large-scale magnetic fields (see Gough \\& McIntyre 1998, for example), and turbulent stresses within another convective region. It is the latter we are mostly interested in studying in this work, namely the case of stars with a convective envelope {\\it and} a convective core.  \\begin{figure}[h] \\centerline{\\epsfig{file=f3.epsf,width=8cm}} \\caption{Illustration of the expected flow pattern for rapidly rotating stars with two convective zones. Gyroscopically pumped flows from one convection zone can enter the radiative zone and return within the second convective zone where the Taylor-Proudman constraint does not apply (see main text for detail). Note the role of the tangent cylinder in such a model in delimiting mixed regions from non-mixed regions. } \\label{fig:gyro2} \\end{figure}   This paper is organized as follows. We first present a fairly exhaustive study of gyroscopic pumping for rapidly rotating stars (in the limit where $\\sigma < 1$). In \\S\\ref{sec:model} we lay out the general formulation of the problem and the assumptions made. Note that the basic model used in this paper is largely inspired from the work of GAA09 but extends it to the case of multiple convective regions and uses a more generally applicable formalism for the fluid dynamics, namely the anelastic approximation. In \\S\\ref{sec:cart} we first solve the problem analytically in a Cartesian coordinate system, using a much simplified stellar model. This exercise provides insight onto how the meridional flow velocities induced by gyroscopic pumping scale with the forcing mechanism (in this case the differential rotation of the inner and/or outer convection zone), with the stratification within the radiative region, and with the system's geometry (i.e. the respective widths of the inner and outer convective regions). In \\S\\ref{sec:spher}, we then apply the same model in a two-dimensional spherical geometry, using a more realistic stellar model as the background state. By comparing Cartesian-model results with equivalent spherical-model results, we deduce a very simple rule to go from one to the other. This rule is particularly useful since one-dimensional Cartesian model numerical solutions can be obtained in a tiny fraction of the time necessary to integrate the two-dimensional spherical case, and can also be pushed to true stellar parameter values (which cannot be done in two dimensions).  Next, we present a simple application of this theory to estimate the surface abundances of lithium (Li) and beryllium (Be) in young Main Sequence stars as a function of their age and mass. More precisely, we are interested in young stars in the mass range of $1.3M_\\odot-1.5M_\\odot$, which have the well-known property of being significantly depleted in Li and Be in their surface layers as compared with slightly more and slightly less massive stars (Boesgaard \\& Tripico 1986; see Boesgaard 2005 and Anthony-Twarog et al. 2009 for reviews). These so-called ``Li-dip'' stars are unique in the sense that they have two significant convection zones (one inner and one outer) which, as described earlier, could promote large-scale mixing between the surface and the interior by gyroscopic pumping. In \\S\\ref{sec:lidip}, we show that the Li and Be depletion rates as induced by gyroscopic pumping for Li-dip stars can be very significant. For reasonable model assumptions, the cool (low-mass) side of the dip is readily explained by gyroscopic pumping, while depletion fractions on the hot (high-mass) side are much larger than observed. When the diffusion of chemical species back into the outer convection zone (by overshooting motions for example) is taken into account, good agreement between the model and the data is achieved on both sides of the dip. Finally, we conclude in \\S\\ref{sec:ccl} by summarizing our main results and discussing future prospects. ",
        "conclusions": "\\label{sec:ccl}  In this paper we performed an exhaustive study of a non-local source of rotational mixing called ``gyroscopic pumping'', which was originally studied in the context of the Earth's atmosphere by Haynes et al. (1991) and later discussed in the case of the Sun by Gough \\& McIntyre (1998), McIntyre (2007), and GAA09.  In this mechanism, large-scale meridional motions are driven by angular-momentum conservation whenever fluid undergoes forces in the azimuthal direction (see \\S\\ref{sec:intromix} and in particular Figure \\ref{fig:gyro} for detail). This is exactly the case in stellar convective zones, where rotationally influenced convection gives rise to the so-called $\\Lambda-$effect (cf. R\\\"udiger 1989; see also Garaud et al. 2010), and typically leads to the azimuthal acceleration of equatorial regions and deceleration of the poles. The resulting large-scale meridional circulation has fluid flowing outward from the rotation axis near the equator, and toward the rotation axis near the pole. Note that this mechanism is quite generic, since it is simply based on angular momentum conservation. However, by contrast with other well-studied significant sources of rotational mixing, it does not rely on stellar spin-down to be effective -- it is an inherently quasi-steady mechanism.  As first shown by GAA09 and studied more extensively here, whether the fluid gyroscopically pumped in the convective zone ends up mixing the nearby radiative zone or not depends on many factors. The case of stars with a single convection zone, in the absence of any other mechanism, was first studied by GAA09. In this case the overall amplitude of the flows penetrating into the radiative zone is limited to slow Ekman flows, which are unlikely to play any significant role in mixing the stellar interior. However, this conclusion could be different if a large-scale magnetic field is present (Gough \\& McIntyre 1998), as it is for instance thought to be the case in the Sun. We defer the magnetic case to a subsequent paper.  Here we focused on stars with two convection zones, in the absence of magnetic fields (we accept that this approximation is probably over-simplistic). We found that gyroscopic pumping provides a significant source of non-local mixing in the radiative zones of these stars, more precisely along the rotation axis, within the cylinder tangent to the inner core (see Figure \\ref{fig:gyrofinal}). This mixing readily explains, for example, the cool side of the well-known Li-dip and, when moderated by the effects of diffusion, provides predictions for the Li- and Be-depletion fractions which give a surprisingly good fit to the data on both sides of the dip (see Figure \\ref{fig:libewithdiff}). It is important to note that contrary to previous models, turbulent mixing here is needed not to explain Li destruction but to explain the replenishment of the surface layers in Li. Of course, the various other effects described in \\S\\ref{sec:liintro}, which were not taken into account here, could also affect the Li and Be abundances: mixing by gravity waves, large-scale rotational mixing induced by stellar spin-down, radiative acceleration, etc. A more sophisticated study of the Li-dip in the light of gyroscopic pumping, taking these effects into account as well as stellar evolution, is deferred to a subsequent publication. Furthermore, it is interesting to note that the gyroscopic pumping mechanism probably plays a role in the global redistribution of angular momentum within the star. Whether this is related to the dichotomy between slow rotators and fast rotators on either side of the dip is an interesting question, which we hope to address in the future.  Gyroscopic pumping, as studied here, also opens up other interesting observational prospects beyond Li-dip stars. For example, another important evolutionary phase when      stars have two convective zones is the core helium burning phase on the red giant branch. Our study shows that this pumping mechanism could provide an important connection between the convective envelope and the core, and perhaps help reconciliate some of the long-standing discrepancies in surface abundances of CNO elements between models and observations. It is also interesting to note that the pumping itself has a tendency to drive a strong azimuthal shear in the system, close to the tangent cylinder (see Figure \\ref{fig:sample}). This shear could perhaps become unstable for very rapidly rotating stars, leading to much more violent mixing in the radiative zone. Alternatively, this shear layer could also become the seat of an $\\Omega-$effect and generate strong toroidal fields around the tangent cylinder. It is interesting to speculate on what kind of field such a flow structure would sustain, and what observable feature it may lead to -- although a full answer to that question will require a study of gyroscopic pumping in the presence of magnetic fields. \\nocite{z92,wa97,tz97,tc03,tc98,sz92,sal93a,sal93b,sal10,sal90,sal98,r89,rm93,r07,p97,mci07,m86,k90,hal91,g07,gmi98,gls91,gg08,2001PhDT,gp08,gaa09,gal10,dp97,cc07,jcdal96,bk02,bt86,b05,b87,bal07,bal96,atal09}"
    },
    "1005/1005.3337_arXiv.txt": {
        "abstract": "{Interstellar dust grains efficiently absorb and scatter UV and optical radiation in galaxies, and therefore can significantly affect the apparent structure of spiral galaxies. We discuss the effect of dust attenuation on the observed structural properties of bulges and discs. We also present some first results on modelling the dust content of edge-on spiral galaxies using both optical and Herschel far-infrared data. Both of these results demonstrate the complex interplay of dust and star light in spiral galaxies. } ",
        "introduction": "In has been known for a long time that interstellar dust grains are an important component of the interstellar medium in galaxies: they efficiently absorb and scatter UV and optical radiation, play an important role in the chemistry of the ISM and are the dominant source of far-infrared and submillimetre emission. Detailed knowledge of the quantity, spatial distribution and physical properties of the dust in spiral galaxies is still controversial in many ways. For several decades after the work by Holmberg (1958), it was assumed that spiral galaxies were optically thin for optical radiation. This conventional viewpoint was questioned in the late 1980s and early 1990s by several teams, including Disney et al.\\ (1989), Valentijn (1990) and Burstein et al.\\ (1991). These authors came to the conclusion that spiral galaxies are optically thick, even in the outer regions. Around the same time, other teams reached completely different conclusions, sometimes on the basis of identical data sets (e.g.\\ Huizinga \\& van Albada 1992).  In retrospective, three issues conspired to complicate the discussion on the optical thickness of spiral galaxies. The first was observational biases in the classical tests, such as the expected variation of isophotal diameter, mean surface brightness with increasing inclination. It is very difficult to quantify these biases, let alone to circumvent them (e.g.\\ Davies et al.\\ 1993). A second issue that complicated the discussion was the unavailability of reliable far-infrared observations that traced the bulk of the dust mass in spiral galaxies. In particular, almost no far-infrared or submillimetre data were available redwards of 100~$\\mu$m, the longest IRAS wavelength. The final and possibly the most important reason why the contrary results were obtained was the simplicity of the models used to analyze or interpret the data. This was convincingly demonstrated by Disney et al.\\ (1989), who showed that simple optically thick models could reproduce the same observations on the basis of which Holmberg (1958) had reached his conclusions.  One of the first to realize that realistic models for dusty galaxies were absolutely needed to investigate the dust content of spiral galaxies was Ken Freeman. He and his Ph.D. student Yong-Ik Byun embarked on the first systematic effort to investigate the complex interplay of dust and star light in realistic models of spiral galaxies, using detailed radiation transfer simulations including both absorption and scattering. The result of this seminal work, presented in Byun et al.\\ (1994), is still a monument in the extragalactic radiative transfer community.  In the past 15 years, the modelling of the interplay between dust and star light in spiral galaxies has been dealt with by many authors.  In particular, several groups have developed the necessary numerical codes to solve the radiative transfer problem in extragalactic environments. The most advanced of these codes are not restricted to only absorption and scattering in 2D geometries, but take into account thermal emission by dust, polarization, kinematics and multi-phase dust distributions (e.g.\\ Popescu et al.\\ 2000; Gordon et al. 2001; Misselt et al.\\ 2001; Steinacker et al.\\ 2003; Bianchi 2008). Our group has also developed a 3D radiative transfer code, SKIRT, based on the Monte Carlo method. Its original aim was to investigate the effects of dust absorption and scattering on the observed kinematics of elliptical galaxies (Baes \\& Dejonghe 2000, 2002; Baes et al.\\ 2000). It has now developed into a mature radiative transfer code that can be used to simulate images, spectral energy distributions, kinematics and temperature maps of dusty systems, ranging from circumstellar discs to AGNs (e.g.\\ Vidal et al.\\ 2007; Stalevski et al.\\ 2010). In particular, the code has been used as the main tool for a detailed investigation of radiative transfer in spiral galaxies (Baes \\& Dejonghe 2001a, b; Baes et al.\\ 2003). In the remainder of this paper we shortly describe two recent results that demonstrate the complex interplay of dust and star light in spiral galaxies: in Section~2 we discuss the effect of dust attenuation on the observed structural properties of bulges and discs, and in Section~3 we present some first results on modelling the dust content of edge-on spiral galaxies using both optical and far-infrared data. ",
        "conclusions": ""
    },
    "1005/1005.3047_arXiv.txt": {
        "abstract": "Open clusters belonging to star-forming complexes are the leftovers from the initial stellar generations. The study of these young systems provides constraints to models of star formation and evolution as well as to the properties of the Galactic disc. We aimed at investigating NGC\\,1981, a young open cluster in the Orion Nebula Region, using near-IR and $BV(RI)_C$ photometric data. We devised a method that accounts for the field contamination and allows to derive photometric membership for the cluster stars. A new cluster centre was determined by Gaussian fittings to the 2-D stellar distribution on the sky, and has been used used to obtain the radial stellar density profile and the structural parameters. Mass functions were computed for stars inside the cluster limiting radius and total mass estimated from them. Although more easily distinguished by its grouping of 6 relatively bright stars, an underlying population of faint pre-main sequence stars is evident in the cluster area. We showed that this population is related to the cluster itself rather than to the nearby Orion Nebula cluster. Additionally a fraction of the cluster low mass stars may have been evaporated from the region in its early evolution leading to the present sparse, loose structure. The estimated parameters of NGC\\,1981 are core radius $R_\\mathrm{c}=0.09\\pm0.04$ pc, limiting radius $R_\\mathrm{lim}=1.21 \\pm 0.11$ pc, age $t=5\\pm1$ Myr, distance modulus ($m-M$)$_0=7.9\\pm0.1$ ($380 \\pm 17$ pc), reddening $E$($B-V$)$=0.07\\pm0.03$ and total mass $m = 137 \\pm 14$ M$_\\odot$. ",
        "introduction": "The open clusters' fundamental physical parameters are important pieces of information for studies on the formation and evolution of the Galactic disc \\citep[e.g., ][]{pca95,Jacobson,Sestito} and as grounding tests for star formation and evolution models \\citep[e.g., ][]{lvd96,d02,Siess2000}. \\citet[hereafter DAML02]{DAML02} have summarized the available information on open clusters in a major catalogue which continuously expands.  In version 2.10\\footnote{http://www.astro.iag.usp.br/$\\sim$wilton} of this catalogue, 1787 open clusters are listed out of which 54 per cent have known distance, age and reddening, 24 per cent have also proper motion and radial velocity and only 10 per cent have  metal abundance determined. \\citet{Kharchenko2005} have created an open cluster catalogue with uniformly determined parameters, among them cluster radius, core radius and age, by means of an automated computational algorithm (The Catalogue of Open Cluster Data, COCD)\\footnote{http://www.univie.ac.at/webda/cocd.html}. This catalogue has been used to investigate the local ($<$\\,1 kpc) population of star cluster complexes revealing new constraints on the Galaxy's structure and kinematics \\citep{Piskunov}.  Sky surveys like 2MASS \\citep{2mass} produced large amounts of near-IR data and have contributed to the discovery of even more objects. Additionally, several studies have been benefited from 2MASS database by employing near-IR photometric analyses of structural and populational properties of Milky Way star clusters \\citep[e.g.,][]{sbb05,Pavani2007,bsb06}. Because the 2MASS database covers the whole sky, it allows data extraction from spatially unlimited regions. Also, near-IR wavelengths are particularly sensitive to discriminate cluster stars from the contaminating field for young stellar systems \\citep[e.g.,][]{sbb05}.  NGC\\,1981 (OCl\\,525, C\\,0532-044) is a young, sparse cluster located $\\sim$1$^{\\circ}$ North from the Orion Nebula at Galactic coordinates $\\ell=208.09^{\\circ}$ and $b=-18.98^{\\circ}$. As noted by \\citet{sha52}, \\emph{the Orion Nebula seems to be separated from NGC\\,1981 only as a result of a somewhat nearer obscuring cloud.} NGC\\,1981 depicts the northern end of Orion's Sword. According to \\citet{sgsb95}, NGC\\,1981 and Collinder\\,70 constitute a possible binary cluster being less than 20 pc apart. The cluster belongs to Gould's Belt, a planar distribution of O and B stars inclined $\\sim$\\,20$^{\\circ}$ with respect to the Galactic plane \\citep[][and references therein]{p01}. In a detailed study of Gould's Belt \\citep{l68}, the brightest stars in the cluster (HD\\,37016, HD\\,37017 and HD\\,37040) had their spectral types determined (B2.5V, B1.5V, B2.5IV, respectively) and photometric distances estimated (all at $494$ pc).  Subsequent works emphasized kinematical properties of NGC\\,1981 stars, among other clusters, aiming at obtaining dynamical and structural characteristics of our Galaxy \\citep[e.g., ][]{h87}. More recently, \\citet{Kharchenko2005} performed an analysis of the cluster as a stellar system deriving reddening $E$($B-V$)$=0.05$, distance modulus ($V-M_v$)$=8.16$ ($d=400$ pc), $\\log{t({\\rm yr})}=7.50$, cluster radius $R=0.25^{\\circ}$ (1.7 pc) and core radius $R_\\mathrm{c}=0.13^{\\circ}$ (0.9 pc). We should note, however, that the age estimated by \\citet{Kharchenko2005} relies upon a single star, as a consequence of the method applied, which progressively removes stars with low membership probabilities from the CMD by means of an iterative process.  Most studies of the cluster are inserted in more general analyses of the Orion star formation complex.  It is part of Orion OB1 association, subgroup c, which is 2 to 6\\,Myr old and located closer to us ($\\approx$400 pc) than the younger Orion Nebula by at least 10 pc \\citep{Bally2008}. However, such a division in subgroups has been questioned in favour of a continuous star forming event \\citep{Muench}. Since NGC\\,1981 bright stars are early B spectral types, a plausible evolutionary sequence would entail supernovae explosions from O-type progenitors causing the compression of the interstellar medium and formation of the present younger populations. On the kinematical side, \\citet{Dias06} derived membership probabilities on the basis of UCAC2 proper motions aided by a statistical method for 158 stars within 15 arcmin of the cluster centre.  \\citet{eac09} investigated hierarchical star formation in Gould's Belt based on the spatial and kinematical distribution of star clusters, either bound star systems or transient stellar condensations with a mean lifetime of $\\sim$\\,10\\,Myr or less. NGC\\,1981 was classified by them as a bound cluster in view of its age ($\\sim$\\,30\\,Myr) as determined by \\citet{Kharchenko2005}. In the present work, we argue  that NGC\\,1981 is indeed younger, as has been shown to be the case for Platais\\,6 and NGC\\,2546, which also belong to Gould's Belt \\citep{eac09}.  Despite the efforts to better understand the astrophysical processes occurring in the Orion Nebula and its surroundings and their connection with Gould's Belt, NGC\\,1981 have not been given particular attention \\citep[except for][]{Kharchenko2005} and, as a consequence, its parameters can be improved. In this work, we present Johnson-Cousins $BV(RI)_C$ photometric observations carried out at the Observat\\'orio do Pico dos Dias (Itajub\\'a, Brazil) along with near-IR photometric data extracted from 2MASS survey to study the open cluster structural properties and stellar content, aiming at a better determination of radius, age, distance, reddening and mass.  This work is structured as follows. Sect. 2 describes the observational data employed. Sect. 3 presents the procedure to obtain the cluster centre, a fundamental step to build a reliable radial density profile, discussed in Sect. 4. Colour-magnitude diagrams are employed in Sect. 5, where a statistical decontamination algorithm is applied to the cluster CMDs, and in Sect. 6, where a comparison of data with isochrones allows a determination of the cluster's astrophysical parameters. In Sect. 7 the radial mass function of the cluster is analysed and in Sect. 8 the stellar density charts used to map the spatial distribution of stars and their frequency in different CMD regions are described. Further discussions and conclusions are given in Sect. 9. ",
        "conclusions": "We investigated the stellar object NGC\\,1981, a young star cluster in the Gould's Belt nearby the Orion Nebula Cluster containing a small population of massive B type stars. We employed near-infrared data from 2MASS to calculate the cluster centre at $\\alpha=5^\\mathrm{h} 35^\\mathrm{m} 08^\\mathrm{s}$, $\\delta=-04^\\circ 20\\arcmin35\\arcsec$ and to subsequently derive the structural parameters from King-profile fittings obtaining a core radius $R_\\mathrm{c}=0.09 \\pm 0.04$ pc and central density $\\sigma_0 = (2\\pm1)\\times 10^2$ stars.pc$^{-2}$ . A limiting radius of $R_\\mathrm{lim}=1.21 \\pm 0.11$ pc was also derived from the radial density profile.  We devised a decontamination procedure based on the method by \\citet{bb07} to statistically remove the underlying field population from the cluster CMD by using an offset-field to sample the background contamination. Tested on NGC\\,1981 with multiple offset-fields, the procedure reliably presented an average field-star exclusion efficiency of 84 per cent, separating cluster members with an average deviation of 5 per cent in the number of stars. It also provides photometric membership for the member stars with average uncertainty of 6 per cent, depending on the selection of the offset-field.  By using optical $BV(RI)_C$ data alongside 2MASS data we performed isochrone fittings on the decontaminated data and determined average values of reddening $E(B-V)=0.07\\pm 0.03$, distance modulus $(m-M)=7.9\\pm0.1$ ($d=380\\pm 17$ pc) and age $5\\pm1$ Myr. The relation between the turn-on $M_V$ and age \\citep{Cignoni2010} was also used as an additional check for the cluster age. In this case the age derived was $3.3\\pm2.7$ Myr. The scarcity of bright stars make this value unreliable as sole indicator of age. We use it to infer that the luminosity of the turn-on is consistent with the age and distance modulus found by isochrone fittings.  By comparing our derived memberships with proper motions from UCAC2 catalogue, we showed that while the small statistical distance between cluster and field populations may have hardened their separation in the VPD, the photometric method was capable to discern these very entangled populations. Indeed, by comparing our derived memberships with proper motions from UCAC3, we were able to distinguish a clear concentration of member stars over a sparse distribution of non-member stars in the VPD.  We derived mass functions for stars inside the inner 5.5 arcmin and 11 arcmin (limiting radius) and evaluated the total mass inside these regions: $m =  30 \\pm 6$ M$_\\odot$ and $m =  107 \\pm 13$ M$_\\odot$, respectively. After fitting a power law to the data, the calculated slopes $\\chi = 0.65\\pm 0.08$ (inner region) and $\\chi = -0.44\\pm 0.03$ (outer region) are flatter than the Salpeter MF, indicating a depletion of low mass stars in the outer region relative to the inner region. Cluster evolution and mass segregation can explain this effect if the cluster has had time to dynamically evolve. In fact, by using proper motion data from DALM02 for stars inside the limiting radius we have derived $t_\\mathrm{cross}=0.15$ Myr and $t_\\mathrm{relax}=0.41$ Myr, meaning that the cluster stars had time to interact gravitationally with each other forming binaries and/or multiple systems composed by massive stars that settle in the cluster centre and leading to the ejection of low mass stars towards the cluster outer regions.  The cluster might have undergone mass loss in the earlier stages of the star forming process leading to the present sparse structure. The energy released into the medium by the evolution of the most massive stars (i.e. by supernovae and stellar winds) may have cleared the cluster from its parental cloud of gas and dust causing a collapse of the core due to the changing gravitational potential and subsequent evaporation of stars in its outer regions.  Through stellar density maps we discriminated the presence of a embedded red population of stars just south of NGC\\,1981. These stars are likely members of the young object NGC\\,1977 and still contaminate the southern population of NGC\\,1981. These maps also demonstrate the power of the procedure employed to disentangle field-stars from the cluster population.  As the majority of young open clusters, NGC\\,1981 is not expected to live longer than a few Myr, evolving from its actual state of marginally bound system to a loose stellar association and finally dispersing itself into the Galactic disc."
    },
    "1005/1005.3271_arXiv.txt": {
        "abstract": "Massive stars influence their surroundings through radiation, winds, and supernova explosions far out of proportion to their small numbers. However, the physical processes that initiate and govern the birth of massive stars remain poorly understood. Two widely discussed models are monolithic collapse of molecular cloud cores and competitive accretion. To learn more about massive star formation, we perform and analyze simulations of the collapse of rotating, massive, cloud cores including radiative heating by both non-ionizing and ionizing radiation using the FLASH adaptive mesh refinement code. These simulations show fragmentation from gravitational instability in the enormously dense accretion flows required to build up massive stars. Secondary stars form rapidly in these flows and accrete mass that would have otherwise been consumed by the massive star in the center, in a process that we term fragmentation-induced starvation. This explains why massive stars are usually found as members of high-order stellar systems that themselves belong to large clusters containing stars of all masses. The radiative heating does not prevent fragmentation, but does lead to a higher Jeans mass, resulting in fewer and more massive stars than would form without the heating. This mechanism reproduces the observed relation between the total stellar  mass in the cluster and the mass of the largest star. It predicts strong clumping and filamentary structure in the center of collapsing cores, as has recently been observed. We speculate that a similar mechanism will act during primordial star formation. ",
        "introduction": "\\label{sec:intro} Understanding the formation of massive stars is of pivotal importance in modern astrophysics. Massive stars are rare and short lived. However, they are also very bright and allow us to reach out to the far ends of the universe. For example, the most distant galaxies in the Hubble Ultra Deep Field \\citep{beckwithetal06} are all characterized by vigorous high-mass star formation. Understanding the origin of massive stars, at present and at early times, therefore is a prerequisite to understanding cosmic history. In our own Milky Way, high-mass stars contribute the bulk of the UV radiation field. This radiation can ionize hydrogen and dissolve molecules. Expanding \\hii\\ regions, bubbles of ionized gas surrounding massive stars, have been identified as an important source  of interstellar turbulence. The complex interplay between gas dynamics and radiation thus is a key element of the evolution of the interstellar medium (ISM) and the Galactic matter cycle. Furthermore, stars are  the primary source of chemical elements heavier than the hydrogen, helium, and lithium that were produced in the Big Bang. Again, massive stars have a disproportionally large share in the enrichment history of the Galaxy and the universe as a whole. It is the violent supernova explosions associated with the death of high-mass stars that contribute most of the metals. Because they insert large amounts of energy and momentum input, supernovae at the same time strongly stir up the ISM and provide for the effective mixing of the newly bread elements.  Despite their importance, the physical processes that initiate and control the build-up of massive stars are still not well understood and subject of intense debate \\citep{maclow04,zinnyork07,mckee07}. Because their formation time is short, of order of $10^5\\,$yr, and because they form deeply embedded in massive cloud cores, very little is known about the initial and environmental conditions of high-mass stellar birth. In general high-mass star forming regions are characterized by more extreme physical conditions than their low-mass counterparts, containing cores of size, mass, and velocity dispersion roughly an order of magnitude larger than those of cores in low-mass regions \\citep[e.g.][]{jijmeyad99,garliz99,kurtzetal00,beutheretal07,motteetal08}. Typical sizes of cluster-forming clumps are $\\sim 1\\,$pc, they have mean densities of $n \\sim 10^5$ cm$^{-3}$, masses of $\\sim 10^3\\,$M$_\\odot$ and above, and velocity dispersions ranging between $1.5$ and $4 \\,$km$\\,$s$^{-1}$. Whenever observed with high resolution, these clumps break up in even denser cores, that are believed to be the immediate precursors of single or gravitationally bound multiple massive protostars.  Massive stars usually form as members of multiple stellar systems \\citep{hohaschik81,lada06,zinnyork07} which themselves are parts of larger clusters \\citep{ladalada03,dewitetal04,testietal97}. This fact adds additional challenges to the interpretation of observational data from high-mass star forming regions as it is difficult to disentangle mutual dynamical interactions from the influence of individual stars \\citep[e.g.][]{gotoetal06,linzetal05}. Furthermore, high-mass stars reach the main sequence while still accreting. Their Kelvin-Helmholtz pre-main sequence contraction time is considerably shorter than their accretion time. Once a star has reached a mass of about $10\\,$M$_\\odot$ its spectrum becomes UV dominated and it begins to ionize its environment. This means that accretion as well as ionizing and non-ionizing radiation needs to be considered in concert \\citep{keto02b,keto03,keto07,petersetal10a,petersetal10b}. It has been realized decades ago that in simple 1-dimensional collapse models the outward radiation force on the accreting material should be significantly stronger than the inward pull of gravity \\citep{larsstarr71,kahn74,wolfcas87} in particular when taking dust opacities into account. Since stars as massive as $100 - 150\\,$M$_\\odot$ have been observed \\citep{bonanosetal04,figer05,rauwetal05} a simple spherically symmetric approach to high-mass star formation is doomed to fail.  Consequently, two different models for massive star formation have been proposed. The first one takes advantage of the fact that high-mass stars always form as members of stellar clusters. If the central density in the cluster is high enough, there is a chance that low-mass protostars collide and so successively build up more massive objects \\citep{bonbatezin98}. As the radii of protostars usually are considerably larger than the radii of main sequence stars in  the same mass range this could be a viable option. However, the stellar densities required to produce very massive stars are still extremely high and seem inconsistent with the observed values of Galactic star clusters \\cite[e.g.][and references therein]{pozwetal10}. An alternative approach is to argue that high-mass stars form just like low-mass stars by accretion of ambient gas that goes through a rotationally supported disk caused by angular momentum conservation \\citep{mckeetan03}. Indeed such disk structures are observed around a number of high-mass protostars \\citep{chini04,chini06,jiangetal08,daviesetal10}. Their presence breaks any spherical symmetry that might have been present in the initial cloud and thus solves the opacity problem. Radiation tends to break out along the polar axis, while matter is transported inwards through parts of the equatorial plane shielded by the disk. Hydrodynamic simulations in two \\citep{yorke02} and three (\\citealt{krumholzetal09}, \\citealt{kuiperetal10}, in prep.) dimensions using a flux-limited diffusion approach to the transport of non-ionizing radiation strongly support this picture. \\citet{krumholzetal09} find non-axisymmetric Rayleigh-Taylor instabilities using gray radiation transfer, whereas \\citet{kuiperetal10}, using frequency-dependent radiation transfer, do not find such instabilities, but nevertheless find strong disk accretion. Strong accretion continues when ionizing radiation is included using ray-tracing methods in three dimensions \\citep{daleetal05,daleetal07b,petersetal10a,petersetal10b}. In a clustered environment or in individual collapsing cores where the disk becomes gravitationally unstable, the three-dimensional models show that material flows along dense, opaque filaments whereas the radiation escapes through optically thin channels of low-density material. It has been demonstrated that even ionized material can be accreted, if the accretion flow is strong enough. \\hii\\ regions are gravitationally trapped at that stage, but soon begin to rapidly fluctuate between trapped and extended states, in agreement with observations \\citep{petersetal10a,galvmadetal10inprep}. Over time, the same ultracompact \\hii\\ region can expand anisotropically, contract again, and take on any of the observed morphological classes \\citep{woodchurch89,kurtzetal94,petersetal10a,petersetal10b}. In their extended phases, expanding \\hii\\  regions drive bipolar neutral outflows characteristic of high-mass star formation \\citep{petersetal10a}.  Another key fact that any theory of massive star formation must account for is the apparent presence of an upper mass limit. No star more massive than  $100 - 150\\,$M$_\\odot$ has been observed \\citep{massey03}. This holds for the Galactic field, however, it is also true for young star clusters that are massive enough so that purely random sampling of the initial mass function (IMF) \\citep{kroupa02,chabrier03} without upper mass limit should have yielded stars above $150\\,$M$_\\odot$ (\\citealt{weidkroup04,figer05,oeycla05,weidetal10}, see however, \\citealt{selmel08}). This immediately raises the question of what is the physical origin of this apparent mass limit. It has been speculated before that radiative stellar feedback might be responsible for this limit \\citep[see, e.g.,][]{zinnyork07} or alternatively that the internal stability limit of stars with non-zero metallicity lies in this mass regime \\citep{appen70a,appen70b,appen87,baraetal01}. However, fragmentation could also limit protostellar mass growth. Indeed, this is what we see in the simulations discussed here. The likelihood of fragmentation to occur and the number of fragments to form depends sensitively on the physical conditions in the star-forming cloud and its initial and environmental parameters \\citep[see, e.g.,][]{krumetal10,girietal10}. Understanding the build-up of massive stars, therefore, requires detailed knowledge about the physical processes that initiate and regulate the formation and dynamical evolution of the molecular clouds these stars form in \\citep{vazsemetal09}.  We argue that ionizing radiation, just like its non-ionizing, lower-energy counterpart, cannot shut off the accretion flow onto massive stars. Instead it is the dynamical processes in the gravitationally unstable accretion flow that inevitably occurs during the collapse of high-mass cloud cores that control the mass growth of individual protostars.  Accretion onto the central star is shut off by the fragmentation of the disk and the formation of  lower-mass companions that intercept inward moving material. We call this process fragmentation-induced starvation and argue that it occurs unavoidably in regions of high-mass star formation where the mass flow onto the disk exceeds the inward transport of matter due to viscosity only and thus renders the disk unstable to fragmentation. We speculate that fragmentation-induced starvation is important not only for present-day star formation but also in the primordial universe during the formation of metal-free Population III stars. Consequently, we expect these stars to be in binary or small number multiple systems and to be of lower mass than usually inferred \\citep{abeletal02,brommetal09}. Indeed, current numerical simulations provide the first hints that this might be the case \\citep{clarketal08,turketal09,stacyetal10}.  In the current study, we analyze the simulations by \\citet{petersetal10a} with special focus on the mass growth history of the individual stars forming and the physical processes that influence their accretion rate. We briefly review the  numerical method we use and the assumptions and approximations behind it in Section \\ref{sec:method}. Then we describe our findings in Section \\ref{sec:results} and discuss them in the context of present-day and primordial star formation in Section \\ref{sec:discussion}. We summarize and conclude in Section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We have presented a detailed analysis of the fragmentation-induced starvation scenario, which was first introduced by \\citet{petersetal10a}. We have studied the accretion history and disk instability in collapse simulations of massive star formation. We have compared the results from a full simulation with radiation feedback and multiple sink particles with two control runs, one with multiple sink particles but without radiation feedback and one with radiation feedback but only a single sink particle. The combined analysis of all three simulations allows us to establish fragmentation-induced starvation as a workable model of massive star formation. We have also compared the new model to the monolithic collapse and competitive accretion models and speculated on the relevance of fragmentation-induced starvation for primordial star formation.  The basic principle of fragmentation-induced starvation is star formation in a rotating gravitationally unstable flow. Gravitational instability in such a flow leads to fragmentation of the accretion flow and the formation of companions around the central massive star that starve the central star of infalling material. In our simulations, star formation occurs within a rotationally flattened, disk-like structure. As more and more material falls onto the disk, star formation proceeds radially outwards, keeping the total disk mass roughly constant. The accretion history of individual stars is tightly connected with the position of these stars relative to the dense filaments and low-density voids in the disk. The fact that most if not all high-mass stars are observed in higher-order multiple stellar systems that themselves belong to more massive star clusters provides strong evidence for the widespread occurrence of fragmentation that forms the basis of our discussion.  We find that accretion heating does not prevent fragmentation of the disk, but leads to a higher local Jeans mass. As a result, fewer stars form with than without radiation feedback. The accretion heating shifts the masses of the most massive stars up, but leaves the average stellar mass almost unaffected. The accretion heating does not change the overall star formation rate of the whole stellar cluster. Feedback by ionizing radiation is unable to stop protostellar growth if the accretion flow is strong enough. However, if the accretion flow weakens due to starvation, an \\hii\\ region can expand and terminate the accretion process. The growing \\hii\\ regions reduce the total star formation rate and do not trigger star formation at the ionization front.  Our model is able to explain the observed morphologies \\citep{petersetal10b} and time variability \\citep{galvmadetal10inprep} of ultracompact \\hii\\ regions. We find that we can consistently reproduce the observed relation between the total mass of the star cluster $\\msinks$ and the maximum stellar mass in the cluster $\\mmax$, $\\mmax \\propto \\msinks^{2 / 3}$. This relation seems to be the general outcome of protostellar interaction in a common cluster environment rather than being a signpost of competitive accretion only, as previously claimed. In fact, the dynamical processes discussed here exhibit exactly the opposite behavior of competitive accretion, rather than run-away accretion onto the most massive star together with the suppression of the growth of lower-mass objects, we see that angular momentum conservation and the presence of lower-mass objects limit the mass growth of massive stars.  Our simulations provide evidence for the rejection of proposals that the observed maximum stellar mass of $\\sim 100\\,$M$_\\odot$ is set by radiative feedback. When disk fragmentation is artificially suppressed (run~A) the central protostar accretes material at very high rate unimpeded by the intense UV radiation it emits without any indications of an upper limit \\cite[see also][]{petersetal10a}. When we permit disk fragmentation to occur, it is the process of fragmentation-induced starvation that prevents the stellar mass to become larger than $\\sim 25\\,$M$_\\odot$ with our choice of initial conditions. We expect more massive, more centrally-condensed, and/or slower rotating  cloud cores to lead to more massive protostars. Indeed, extensive parameter studies \\citep{girietal10} show that the initial density profile dominates the accretion behavior, explaining the formation of a $40\\,$M$_\\odot$ star from a $100\\,$M$_\\odot$ core in \\citet{krumholzetal09}. The alternative view is to attribute the apparent stellar mass limit to internal stability constraints."
    },
    "1005/1005.3792_arXiv.txt": {
        "abstract": "\\noindent Spectra of the brightest blazars detected by the {\\em Fermi Gamma-ray Space Telescope} Large Area Telescope cannot be described by a simple power law model. A much better description is obtained with a broken power law, with the break energies of a few GeV. We show here that the sharpness and the position of the breaks can be well reproduced by absorption of $\\gamma$-rays via photon--photon pair production on  He\\,{\\sc ii}  Lyman recombination continuum and lines. This implies that the blazar zone lies inside the region of the highest ionization of the broad-line region (BLR) within a light-year from a super-massive black hole. The observations of $\\gamma$-ray spectral breaks open a way of studying the BLR photon field in the extreme-UV/soft X-rays, which are otherwise hidden from our view. ",
        "introduction": "The {\\em Fermi Gamma-ray Space Telescope} has detected more than a hundred of blazars in the 100 MeV--100 GeV range \\citep{Abdo09_AGN} and allowed their detailed spectral studies with unprecedented accuracy. A surprising result is that the spectra of high-luminosity sources, flat-spectra radio quasars (FSRQs) and low-energy synchrotron peaked BL Lac objects, are much better described by a broken power law than by a simple powerlaw or any smoothly curved models \\citep{Abdo09_3C454.3,Abdo10_blazars}. The break energies, mostly lying in the 2--10 GeV range (as measured in the object frame),  seem to be too small to be produced by the  $\\gamma$-ray absorption due to photon--photon pair production in the broad-line region (BLR), as the strongest BLR line, Ly$\\alpha$, absorbs $\\gamma$-rays starting from only 25.6 GeV.  The observed spectral breaks are also too sharp and the change in the spectral indices are too large to be associated with the cooling or the Klein-Nishina effects \\citep{GT09}.  Calculations of the $\\gamma$-ray absorption within the BLR are usually  limited to the contribution from lines in the observed optical band (e.g. \\citealt{Liu06,Reimer07}) with a few exceptions \\citep{TM09}. Theoretical models, however, predict strong helium UV lines as well as metal lines in the soft X-rays \\citep[see, e.g.][]{Krolik99}, because a typical quasar spectrum has a power law tail extending to the X-ray band \\citep{Laor97}. In addition to the lines, there are strong recombination continua of hydrogen and He\\,{\\sc ii}  at 13.6 and 54.4 eV, respectively. These sharp,  line-like features cause jumps in the $\\gamma$-ray opacity at $\\sim$19.2 and 4.8 GeV. We propose here that the observed spectral breaks in blazars are produced by absorption on photons of these recombination continua.  \\begin{figure*} \\centerline{\\epsfig{file= fig1a.eps,width=6.7cm}\\hspace{2cm}\\epsfig{file= fig1b.eps,width=6.7cm}} \\caption{ (a) Spectrum of the broad emission line region. The spectra for various  ionization parameters $\\xi$ computed using a photoionization code {\\sc xstar} version 2.2 \\citep{KB01} are averaged on a logarithmic energy grid with the 5\\% spectral resolution to show the most significant features dominating the total flux. The incident typical quasar spectrum \\citep{Laor97}  shown by a smooth curve consists of a multicolor disk with the maximum temperature $10^5$ K and a power law tail of photon index $\\Gamma=2$ extending to 100 keV with 10\\% of the total luminosity. (b) Average cross section (in units of Thomson cross section) for $\\gamma$-rays due to photon--photon pair production on the BLR spectra shown in panel (a). Opacity produced by the disk radiation, direct or scattered outside the BLR clouds \\citep[see e.g.][]{BL95}, is neglected here. Jumps in opacity are clearly seen at energies corresponding to the strongest lines and recombination continua. The 0.3--0.5 GeV jump for large $\\xi$ is produced by O\\,{\\sc viii} $\\lambda\\lambda$16--19  lines, the jump at low $\\xi$ is due to the O\\,{\\sc vii}   $\\lambda$22 line complex (see Table \\ref{tbl:blr}). The jump at $\\sim$5 GeV visible at all ionizations is due to the He\\,{\\sc ii} recombination continuum and Ly lines at 40--60 eV. The 20--30 GeV jump seen at low $\\xi$ is produced by H\\,{\\sc i} Ly lines and continuum and the C\\,{\\sc iv} $\\lambda$1549 line. } \\label{fig:blr} \\end{figure*} ",
        "conclusions": "The GeV breaks observed in blazars are well described by $\\gamma$-ray absorption via photon--photon pair production on He\\,{\\sc ii} and H\\,{\\sc i} recombination continuum photons. In RGB J0920+446, the absorption is seen only at high energies with the break energy $E_{\\rm break}$ corresponding to the pair-production threshold on hydrogen recombination photons, while in other cases the break is close to the threshold for the absorption on He\\,{\\sc ii}  recombination continuum. The exact position of the break depends on the ionization parameter that determines the contribution of metals and affects the position of the centroid of the 50 eV complex. If $\\tau_{\\rm He}$ is small, then the break shifts to 19 GeV as observed in RGB J0920+446.   A rather large ratio of the fitted optical depths  $\\tau_{\\rm He}/\\tau_{\\rm H}\\sim 1/4$, implies that the $\\gamma$-ray emitting region has to lie within the high-ionization zone of the BLR with $\\log \\xi>2$. For the brightest $\\gamma$-ray object in our sample, 3C 454.3, with its accretion luminosity of about $10^{47}$ erg s$^{-1}$ \\citep{SEB88}, the high-ionization zone should be within about 0.1 pc. This corresponds to about $10^3$ Schwarzschild radii for an $\\sim10^9 M_{\\odot}$ central black hole \\citep{Bonnoli10}. At such a distance, the  luminosity in the 50 eV complex can be as small as  $\\sim$10$^{44}$ erg s$^{-1}$, a per mille of the accretion luminosity, to provide the necessary opacity with $\\tau_{\\rm He}\\sim 6$ (see Equation \\ref{eq:taut}).   The opacity measured above 20 GeV and the  Ly$\\alpha$ luminosity of 10$^{45}$ erg s$^{-1}$ observed in 3C 454.3 \\citep{Wills95} allow us to estimate the Ly$\\alpha$ emission zone size. Taking the recombination continuum luminosity equal to that of Ly$\\alpha$ and using Equation (\\ref{eq:taut}) we get $R_{\\rm Ly\\alpha}=2\\pm1$ pc, which indicates  that absorption at these energies happens at a larger distance than absorption by   He\\,{\\sc ii} photons.   The constraints obtained on the $\\gamma$-ray emission site imply that the jet is already accelerated to a relativistic velocity within a thousand gravitational radii. It also strongly  constrains the  mechanisms for $\\gamma$-ray production. The possible sources of soft photons for Comptonization in the jet are the accretion disk \\citep{DS93} and the BLR \\citep{SBR94}. As a source of photons, the dust emission at 10 pc scale \\citep{BSM00,SMM08} cannot be important.  Let us also remark that the GeV photons absorbed in the BLR produce electron--positron pairs which, spiraling in the magnetic field, radiate away their energy isotropically. This emission cannot compete with the beamed emission from the jet, but can contribute to the high-energy emission of radio galaxies observed at large angles to the jet axis \\citep[see, e.g.][]{RB10}.   Our interpretation of the GeV breaks implies that additional breaks (depressions) at 0.3--0.7 GeV produced by the soft X-ray lines within the high-ionization zone should be seen, once the  photon statistics is high  enough. The $\\gamma$-ray spectroscopy can be used as a powerful tool for studying the extreme-UV and soft X-ray emission in the quasars' vicinity, which is otherwise hidden from us by interstellar absorption."
    },
    "1005/1005.4384_arXiv.txt": {
        "abstract": "{The young cluster NGC 2264 was observed with the Corot satellite for 23 days uninterruptedly in March 2008 with unprecedent photometric accuracy. We present here the first results of the analysis of the accreting population that belongs to the cluster and was observed by Corot. } {We intended to look for possible light curve variability of the same nature as that observed in the classical T Tauri star AA Tau, which was attributed to a magnetically controlled inner disk warp. The inner warp dynamics is directly associated with the interaction between the stellar magnetic field and the inner disk region. } {We analysed the Corot light curves of 83 previously known classical T Tauri stars that belong to NGC 2264 and classified them according to their morphology. We also studied the Corot light curve morphology as a function of a \\emph{Spitzer}-based classification of the star-disk systems. } {The classification derived on the basis of the Corot light curve morphology agrees very well with the \\emph{Spitzer} IRAC-based classification of the systems. The percentage of AA Tau-like light curves decreases as the inner disk dissipates, from 40\\% $\\pm$ 10\\% in systems with thick inner disks to 36\\% $\\pm$ 16\\% in systems with anemic disks and none in naked photosphere systems. Indeed, 91\\% $\\pm$ 29\\% of the CTTS with naked photospheres exhibit pure spot-like variability, while only 18\\% $\\pm$ 7\\% of the thick disk systems do so, presumably those seen at low inclination and thus free of variable obscuration. } { AA Tau-like light curves are found to be fairly common, with a frequency of at least $\\sim$ 30 to 40\\% in young stars with inner dusty disks. The temporal evolution of the light curves indicates that the structure of  the inner disk warp, located close to the corotation radius and responsible for the obscuration episodes, varies over a timescale of a few ($\\sim$1-3) rotational periods. This probably reflects the highly dynamical nature of the star-disk magnetospheric interaction. } ",
        "introduction": "\\hspace{1.5em}  T Tauri stars are young, optically visible, low-mass stars still contracting toward the main sequence. They have strong magnetic fields ($\\sim$2 kG) and are X-ray emitters. The so-called weak line T Tauri stars (WTTSs) do not exhibit evidence of disk accretion, while the classical T Tauri stars (CTTSs) do. CTTSs present broad emission lines and sometimes also forbidden emission lines. They are spectroscopically and photometrically variable, and show ultraviolet (UV), optical and infrared (IR) excess with respect to the photospheric flux.  Magnetospheric accretion models are the current consensus to explain the main observed characteristics of CTTSs \\citep{shu94,har94,muz01,kur06}. In these models, the stellar magnetosphere is strong enough to disrupt the circumstellar disk at a few stellar radii. Material in the inner disk, ionized by stellar radiation, falls towards the star following magnetic field lines and hits the stellar photosphere at near free-fall velocities, creating hot spots on the stellar surface. Part of the ionized material in the inner disk region is ejected in a magnetically controlled wind. In this scenario, the broad permitted emission lines are formed partly in the accretion funnel and the hot spot emits a continuum flux that is responsible for the UV and optical excess that veils the photospheric lines. The IR excess comes from reprocessing by the dust in the disk of the stellar and accretion radiation and, at least for high accretion rate systems, viscous heating may also contribute to it. When observed, forbidden emission lines are thought to be formed in the low density wind.  In the last decade, numerical simulations of accreting young stars have predicted a very dynamical star-disk interaction, mediated by the stellar magnetic field \\citep{goo99,mat02,rom09,zan09}. Many magneto-hydrodynamical (MHD) model predictions derive from the idea that the stellar magnetic field interacts with the inner disk region near the co-rotation radius but not only at the co-rotation point. Consequently, due to differential rotation between the star and the inner disk region, the magnetic field lines become distorted after a few rotational periods and eventually reconnect, restoring the initial field configuration. This process goes on as the star rotates. The time scale for the reconnection events depends on the diffusivity of the stellar magnetic field lines in the inner disk region, which is a very poorly constrained parameter.  AA Tau is one of a few CTTSs studied with enough detail to test the MHD model dynamical predictions. The star was observed photometrically for a month during three different campaigns, two of which included high-resolution simultaneous spectroscopy \\citep{bou95,bou03,bou07}. AA Tau shows a light curve (LC) with a flat maximum interrupted by deep quasi-periodical minima that vary in depth and width from one rotational cycle to the other. The minima occur with little color change and are thought to be due to obscuration of the stellar photosphere by circumstellar disk material present in an inner disk warp. The warp is due to the interaction of the stellar magnetic field, inclined with respect to the rotation axis, and the inner disk region. The spectroscopic results showed that the magnetic field lines inflate in the course of a few rotational cycles, as measured by the radial velocity of the red and blue absorption components of the \\hal \\ line. Moreover, when the field lines are the most inflated, accretion to the star is suppressed, with no apparent veiling, and low emission-line equivalent widths, which again corroborates the MHD model predictions. After reconnection, accretion starts all over again. The last two AA Tau observational campaigns were separated by 5 years and the field line inflation discovered in the second campaign was still present with the same characteristics in the third one.  Although the AA Tau observations gave strong support to the MHD results, it was still unclear whether this behaviour was common among CTTSs. In order to test this, one needs good photometric measurements of a large number of CTTSs, covering several rotational cycles. In Taurus the typical rotational period of a CTTS is around 8 days, so that stars need to be monitored for at least a month continuoulsy. This is not an easy task from the ground, due to telescope time allocation and weather conditions. The Corot satellite additional program to observe the young stellar cluster NGC 2264 has allowed us to perform such study on a large sample of CTTSs.  NGC 2264 is a well studied young stellar cluster with an age of $\\approx 3$ Myr located at a distance of about 760 pc \\citep[see][for a recent review on the cluster]{dah08}. We matched the various available datasets on the cluster from the literature based on spatial coincidence using a 2'' tolerance. In a small number of cases, double identifications (i.e. more than one object within 2'') were transformed into single ones, specifically in cases in which one of the two objects had a small offset and the other a large one. Due to the many studies on the cluster, going from UV to X-rays, it was possible to establish reasonable criteria for cluster membership. We considered as likely members of NGC 2264 stars selected according to one or more of the following criteria: i) photometric \\hal \\ and variability with the data of \\citet{lam04} and following their criteria, ii) X-ray detection \\citep{ram04, fla06} and location on the cluster sequence in the ($I$, $R-I$) diagram if $R$ and $I$ magnitudes are available, iii) spectroscopic \\hal \\ equivalent width greater than 10 \\AA \\ and iv) \\hal \\ emission line width at 10\\% intensity greater than 270 \\kms, as proposed by \\citet{whi03} to identify accreting T Tauri stars, and measured by \\citet{fur06} for many cluster members.  The mean rotational period of CTTSs in NGC 2264 is around 3 to 4 days (see Sect. \\ref{discussion}). Corot observed the cluster for 23 days uninterruptedly, covering several rotational cycles for most of the stars, thus allowing the identification of AA Tau-like candidates and the determination of precise rotational periods for cluster members. The complete rotation analysis will be discussed separately in another paper (Affer et al. in preparation). ",
        "conclusions": "We showed that the AA Tau photometric behavior is common among CTTSs, being present in 28\\% $\\pm$ 6\\% of the CTTSs in our sample. This represents a lower limit, since the AA Tau-like LCs are more likely produced at high inclinations, and we are probably missing about 20\\% to 30\\% of the very high inclination systems, according to  the calculations by \\citet{ber00}, which will be totally obscured by a flared disk and thus too faint to be observed by Corot.  If our interpretation of such systems is correct, the photometric minima are due to obscuring material located in the inner disk region, near the corotation radius. This material is built up through the interaction between an inclined stellar magnetosphere and the inner disk region. The observed periodical changes in width and depth of the observed minima, over a timescale of a few ($\\sim$ 1-3) rotational periods, would then reflect the dynamics of such an interaction, as predicted by MHD models of young low mass star-disk systems.  We compared the Corot light curves with \\emph{Spitzer} IRAC data of the same systems and showed that the agreement between classifications based on the two datasets is excellent. The percentage of AA Tau-like light curves, which are due to obscuration by circumstellar material in the inner disk region, varies as the inner disk dissipates, decreasing from 40\\%  $\\pm$ 10\\% in systems with thick inner disks to 36\\% $\\pm$ 16\\% in systems with anemic disks and none in naked photosphere systems. Indeed, 91\\% $\\pm$ 29\\% of the systems with naked photospheres exhibit pure spot-like variability, while only 18\\% $\\pm$ 7\\% of the thick disk systems do so, presumably those seen at low inclination and thus free of variable obscuration.  \\hspace{1.5em}"
    },
    "1005/1005.0387_arXiv.txt": {
        "abstract": "We show the advantages of a wedding cake design for Sunyaev-Zel'dovich cluster surveys. We show that by dividing up a cluster survey into a wide and a deep survey, one can essentially recover the cosmological information that would be diluted in a single survey of the same duration due to the uncertainties in our understanding of cluster physics. The parameter degeneracy directions of the deep and wide surveys are slightly different, and combining them breaks these degeneracies effectively. A variable depth survey with a few thousand clusters is as effective at constraining cosmological parameters as a single depth survey with a much larger cluster sample. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.2457_arXiv.txt": {
        "abstract": "This review presents basic equations for the solution of the NLTE radiative transfer problem for trace elements and methods for its solution are summarized. The importance of frequency coupling in radiative transfer in stellar atmospheres is emphasized. ",
        "introduction": "A stellar atmosphere is the transition between the dense optically thick stellar core and the transparent interstellar medium. It is the boundary layer of a star and it is the only part of the star we can see directly. Knowledge of the stellar atmosphere is a gate to the investigation of the stellar interior, and, consequently, to stellar evolution processes. The radiation passing through it serves not only as an information tool describing what is happening there, but it also significantly modifies the stellar atmosphere itself. This makes the task of determining the emergent radiation from the star highly nonlinear and therefore also complicated.  Since the stellar atmosphere is a transition from the optically thick stellar interior to optically thin interstellar medium, the corresponding radiative transfer equation for the whole atmosphere cannot be subject to global simplifications. For the optically thick part of the atmosphere the diffusion approximation is valid. Radiation propagates forward slowly resembling the process of diffusion. This approximation greatly simplifies the radiation transfer equation. However, the diffusion approximation is unusable for the optically thin part of the atmosphere. The situation there is again simple from another point of view. Since the medium is optically thin, very little interaction with matter occurs there, which often allows us to neglect absorption and to take into account only emission. Between these two relatively simple regions there is a transitional region, where neither diffusion nor optically thin approximations are valid. We have to solve the full set of radiative transfer equations together with many constraint equations, which introduce the dependence of the absorption and emission coefficients on radiation and which make the problem to be solved highly nonlinear.  Radiation carries information about the medium where it was created. In stellar atmospheres, it has an additional role. Since the photon mean free path is much larger than the particle mean free path, radiation in stellar atmospheres influences the matter far from the place where it is formed, sometimes very significantly. It changes the population numbers of particular atomic energy levels, sometimes very far from the equilibrium values. These changes then influence spectral lines, which are observed by us. Radiation also transfers energy causing heating (or cooling) of specific parts of the stellar atmosphere. In some cases radiation has such an enormous effect on the atmospheric matter that it becomes the basic reason for a stellar wind, which is then called as 'radiatively driven'. All these crucial effects of radiation influence the transfer of radiation, which becomes strongly nonlinear. ",
        "conclusions": "In this introductory part we emphasized the importance of frequency coupling of radiation (in addition to the spatial and angle coupling) for the conditions of stellar atmospheres. This frequency coupling is ensured mainly by equations of statistical equilibrium, which connect distant parts of the frequency spectrum. We also briefly summarized underlying equations for the NLTE radiative transfer problem in stellar atmospheres with a focus on the problem of its solution for the case of trace elements.  There are two basic methods of the solution of the RTE+ESE system, namely the complete linearization method, which is used by the Kiel code \\citep[see][\\citetalias{Kielcode}]{Kielcode}, and the accelerated lambda iteration method used by the DETAIL code \\citep[see][\\citetalias{DETAILcode}]{DETAILcode}."
    },
    "1005/1005.1809_arXiv.txt": {
        "abstract": "\\vskip0.25cm In order to understand how locally static configurations around gravitationally bound bodies can be embedded in an expanding universe, we investigate the solutions of general relativity describing a space-time whose spatial sections have the topology of a 3-sphere with two identical masses at the poles. We show that Israel junction conditions imply that two spherically symmetric static regions around the masses cannot be glued together. If one is interested in an exterior solution, this prevents the geometry around the masses to be of the Schwarzschild type and leads to the introduction of a cosmological constant.  The study of the extension of the Kottler space-time shows that there exists a non-static solution consisting of two static regions surrounding the masses that match a Kantowski-Sachs expanding region on the cosmological horizon. The comparison with a Swiss-Cheese construction is also discussed. ",
        "introduction": "A key issue as regards the nature of the universe is how the smooth expanding large-scale universe is constructed from small-scale extremely inhomogeneous domains that are essentially static. The space-time in the Solar system for example is static to a very good approximation, but it is very inhomogeneous ($\\delta\\rho/\\rho \\simeq 10^{+30}$). How can one construct an expanding universe solution by gluing together large numbers of such quasi-static domains? The static space-time regions representing local masses and the vacuum regions around stars are of course structured into galaxies, clusters of galaxies, and large scale structures such as walls and voids; the same essential problem arises if we consider how they in turn can be assembled into an expanding universe. For simplicity and analytic clarity, we will assume that: \\begin{enumerate} \\item the component massive entities are exactly spherically symmetric; \\item these entities, whether we think of them as stars, galaxies, or galaxy clusters, are embedded in an exact vacuum solution of the Einstein field equations (possibly with a non-zero cosmological constant, $\\Lambda$); \\item they are far enough apart from each other that the vacuum space-time around each of them is also locally spherically symmetric. \\end{enumerate}  Given the assumption of local spherical symmetry around each mass, it is crucial that Birkhoff's theorem \\cite{Bir23,HawEll73} will apply:\\\\  4. the vacuum solution around each mass will necessarily be locally static.\\\\  \\noindent Thus we do not have to assume local staticity which follows from the other previous assumptions. The local exterior space-time (outside the massive entities) will thus locally be either a Schwarzschild solution (if $\\Lambda = 0$) or a Kottler solution (if $\\Lambda \\neq 0$). So we can ask: how can one join exact Schwarzschild or Kottler solutions together to form an expanding universe?  \\subsection{The Lindquist-Wheeler (L-W) solution} This issue was tackled in a very innovative paper by Lindquist and Wheeler a long time ago~\\cite{LinWhe57}, using a Schwarzschild cell method to model an expanding universe with closed spatial sections (having the topology of a 3-sphere, $S^3$). For simplicity, they used a regular lattice, which allows a very limited set of possibilities. Considering $N$ vertices, they state that ``every vertex can be equidistant from its neighbours only when $N = 5, 8, 16, 120,$ or $600$\", which corresponds to the most homogeneous topologies of the 3-sphere~\\cite{topoS3}. They then derived equations of motion for the expanding universe from junction conditions between the cells. This work has been recently revisited in an interesting way in Refs.~\\cite{CliFer09,CliFer09a}.  This approach is quite different conceptually from the commonly used ``Swiss-cheese\" construction~\\cite{EinStr45,EinStr45a}. In that case, one \\emph{starts off} with a Friedmann-Lema\\^{\\i}tre spatially homogeneous and isotropic geometry, and then cuts out spherical ``vacuoles\" within which individual masses are embedded. These masses are thus contained in vacua within a spatially homogeneous fluid-filled cosmos. In the Lindquist-Wheeler approach on the other hand, one starts off with the inhomogeneous vacuoles alone, and then glues them together to construct an emergent Friedmann-Lema\\^{\\i}tre geometry when averaged on large scales. There is no fluid filling the space-time; rather (as in the case of kinetic theory) fluid-like behaviour emerges on large scales when one coarse-grains over the detailed structure. This is a far more fundamental approach to the study of the relation between locally static inhomogeneity and a globally expanding universe. However, this method is not strictly self-consistent in that the gravitational fields of the neighboring particles would in fact deform the field in the neighborhood of each cell's vertex, thereby resulting in an approximate rather than exact spherically symmetric space-time region (the real solution will be locally a bit anisotropic about each vertex). Nevertheless it is an acceptable approximation, and is a very useful approach to tackle the issue raised here.  \\subsection{Two-mass L-W solutions}  In order to understand the situation better, we investigate the simplest possible such solution with isolated spherically symmetric massive objects embedded in a vacuum with $S^3$ spatial topology: namely the case with $N=2$, which Lindquist and Wheeler did not consider~\\cite{LinWhe57}. With the two masses at antipodal points of the spatial sections, the space-time remains exactly spherically symmetric about each mass - the approximation comment made above no longer applies and an exact solution to the problem can be found. The exact spherical symmetry around each mass ensures that Birkhoff's theorem applies and the space-time is locally static. Setting off radially from either one of the masses, in order to have an $S^3$ spatial topology, the area of the surrounding 2-spheres $S^2(r)$ at proper distance $r$ from the centre of symmetry must reach a maximum and then decrease to zero at the antipodal point: indeed we would expect to have two identical static solutions (if the 2 masses are equal) back to back, joined at the equator of the $S^3$-spatial section.  Two striking results now emerge (as will be shown in details below). Firstly, if we want an exterior solution of this kind, we need a positive cosmological constant; $\\Lambda > 0$. Hence the two local vacuum solutions to be joined together are Kottler space-times \\cite{Kot18,Wey19} rather than Schwarzschild solutions, filled in by regular static bodies (e.g. stars) at their centers, so there is no horizon near the centre of symmetry. Let us call this a Kottler exterior solution.  Secondly, we first show that Israel junction conditions \\cite{Isr66,Isr67} impose that the two space-times can only be joined on an horizon or a surface of maximum area, where one of the metric components vanishes. However, after constructing a coordinate system regular on the horizon, we show that the matching cannot be performed on the horizon. Gluing the two pieces together on a surface of maximum area back to back would result in a static configuration, but this is also shown to be impossible as no coordinate system can regularly cover the neighbourhood of the surface of maximum area. This prevents the existence of a static space-time analog to the Einstein-Static % model with the mass concentrated into two antipodal compact objects. Instead, we find that the only solution contains two static regions matched across their null horizons to a pair of Kantowski-Sachs expanding and contracting universes; the expansion of the universe is possible because of the existence of these spatially homogeneous but anisotropic expanding regions. This solution is obtained by considering the global structure of the exterior Kottler space-time, which we show is essentially the same as that of the maximal de Sitter hyperboloid \\cite{HawEll73,Sch56}.  Surprisingly, this is again rather like the Swiss-Cheese models, but now with the static domains embedded across null horizons (instead of space-like hypersurfaces in the case of the Swiss-Cheese). Each static and each expanding domain covers only part of the resulting maximal Kottler exterior two-mass solution, in a way exactly analogous to what occurs in the maximally extended de Sitter solution. Furthermore, this shows that a globally static solution can exist only at the price of introducing a (unphysical) surface layer.  A globally non-singular 2-mass LW space-time does indeed exist, but it has horizons separating spatially inhomogeneous static domains centred on the two masses from spatially homogeneous time-evolving region. It is this expanding universe domain that allows the two masses to move away from each other, and so is the reason this universe model can expand despite the static nature of space-time near each embedded massive object.  The rest of the text will establish these results and is organized as follows: section II is a quick summary of the statements of Birkhoff's theorem; section III will present the results obtained when one tries to glue together two spherically symmetric static space-times; section IV will discuss the results in a Swiss-cheese configuration and comment on the analogy with the previous gluing approach; finally section V will be a discussion of the results and their implications. ",
        "conclusions": "The LW models are significant tools in elucidating the ongoing discussion on averaging in cosmology, and the relation between small scale inhomogeneous models and large scale spatially homogeneous models (see e.g. \\cite{CliFer09, CliFer09a,Wil09, Wil09a, Maretal07}). This could provide a better understanding of the emergence of a fluid-like description of the Universe on large scales from a more realistic local clumpy distribution of matter. In the line of this program, we have analyzed the geometrical structure of an exact two-mass singularity free version of the LW models: an expanding model which is the analogue of the maximal de Sitter hyperboloid once it has been embedded in a 5 dimensional space-time. This resolves the paradox raised in the introduction to this paper: how can locally static domains be glued together to give an expanding universe? Indeed, the solution presented here is locally static around each compact object but not globally so even if the conditions of the Birkhoff theorem still holds (spherical symmetry around each mass and vacuum + cosmological constant solution of Einstein equations). This is different from a de Sitter space-time, which is locally static everywhere: nothing intrinsic to the de Sitter space-time identifies any particular null surface as the horizon, and there are local timelike and null Killing vector fields at every point. The 2-mass exterior Kottler solution, on the other hand, is not locally static in the region across the horizon: as in the case of the Schwarzschild solution (inside its horizon), there is no timelike Killing vector field there, and the horizon is uniquely defined locally by existence of a null Killing vector field.  We have also shown that it is not possible to construct a globally static solution that is the analogue of the Einstein Static universe (a timelike Killing vector field occurs at every point), but with the mass concentrated in two antipodal massive bodies. This would require the introduction of an unphysical surface layer. We have shown that the expanding region can either be filled with a cosmological constant and has a Kantowski-Sachs geometry, or by a pressureless fluid and a cosmological constant. In that latter case, the geometry of the expanding region can be of the FL type.\\\\  \\noindent{\\bf Acknowledgements:} We thank Timothy Clifton for his fruitful comments on an early version of this text. We also thank Chris Clarkson, Gilles Esposito-Far\\`ese, Charles Hellaby, Jeff Murugan, Tony Rothman and Amanda Weltman for discussions. JPU thanks UCT for hospitality during the period this project was initiated. JL is supported by the Claude Leon Foundation (South Africa)."
    },
    "1005/1005.1512_arXiv.txt": {
        "abstract": "{The Tully--Fisher relation represents a connection between fundamental galaxy parameters, such as its total mass and the mass locked in stars. Therefore, the study of the evolution of this relation in the optical and infrared bands can provide valuable information about the evolution of the individual galaxies through the changes found in each band.} {This work aims to study the Tully--Fisher relation at high redshift in the B, V, R, I, and K$_{\\rm S}$--bands by comparison with the local relations derived from a large sample of galaxies in the redshift range 0.1$<$z$<$0.3, processed in the same way, and with the same instrumental constraints that the high--redshift sample.} {Using the large photometric information available in the AEGIS database, we have studied the best procedure to obtain reliable k--corrections. Instrumental magnitudes are then k-- and extinction corrected and the absolute magnitudes derived, using the concordance cosmological model. The rotational velocities have been obtained from the widths of optical lines using DEEP2 spectra. At high redshift, this method is found to give better results than using the rotation curve, due to spatial resolution limitations. Finally, morphology has been determined via visual classification of the HST images. From the above information, the Tully--Fisher relations in B, V, R, I and K$_{\\rm S}$--bands are derived for the local and high--redshift sample.} {We detect evolution in the B, V and R--bands in the sense that galaxies were brighter in the past for the same rotation velocity. The change in luminosity is more noticeable in the bluer bands. This colour evolution, unnoticed in our previous work (Fern\\'andez Lorenzo et al. 2009) has been detected thanks to the more reliable k--corrections carried out in this paper, which included photometry from B to IRAC bands. The change in the (V--K$_{\\rm S}$) and (R--I) colours (for a fixed velocity) could be interpreted as an ageing of the stellar populations as consequence of the star formation decrease since z=1.25. In addition, we conclude that spiral galaxies could have doubled their stellar masses in the last 8.6 Gyr.} {} ",
        "introduction": "The relation between luminosity and maximum rotational velocity of spiral galaxies \\citep{1977A&A....54..661T} is an important distance estimator, and it has been used for measuring the Hubble constant H$_0$ \\citep[e.g.][]{2000ApJ..533..744}. The study of this relation at different cosmic times can probe evolution in galaxy properties such as its total mass, or the relation between dark and luminous matter \\citep{1995MNRAS.273L..35Z}. Moreover, any consistent model of galaxy formation and evolution should be able to reproduce the evolution of the Tully--Fisher relation (hereafter TFR). Then, its study has important implications in determining fundamental cosmological parameters, in the study of structure formation, and in the evolution of disc galaxies \\citep{2000ApJ...538..477N}.  Traditionally, the internal kinematics of nearby galaxies is measured using 21cm line widths, but the sensitivity of radio telescopes hampers its use at high redshift. Then, optical lines become important in evolutionary studies of TFR, and some authors have studied the relation between both velocity indicators. For example, \\citet{1992ApJS...81..413M} compared the projected rotation velocity measured from H${\\rm \\alpha}$ rotation curves, with the velocity measured from integrated HI profiles at the 50\\% of the maximun flux level. They obtain a difference of 10 km/s, which is attributed to the fact that HI widths measure not only the rotational velocity but the internal galaxy turbulence as well. This turbulence is important mainly in the most external region of the galaxy, where the gravity is lower. However, the optical emission does not extent so far. Then, the contribution of the turbulence to the rotational velocity is negligible when using these lines. \\citet{1996vogt,1997vogt} modeled rotation curves via optical long slit observations, and they found a modest luminosity evolution in the B--band TFR ($\\Delta$M$_B$(high z$-$local)$\\leq$$-$0.4 at $<$z$>$=0.5). However, \\citet{1997MNRAS.285..779R} and \\citet{1998simard} found a stronger evolution ($\\Delta$M$_B$$\\approx$$-$1.5 at z=0.4), modeling the kinematics of disc galaxies in a similar way. Subsequent works have measured the evolution of the TFR, again finding different and even contradictory results. \\citet{2002ApJ...564L..69Z}, and \\citet{2004A&A...420...97B}, find a relation at high redshift (z$=$1) shallower than that measured in local samples, and evidence for a luminosity evolution with look--back time of ${\\Delta}M_B \\approx -1$ magnitudes at redshift z$=$1. In their study, contrary to previous findings, they claim that less massive galaxies were brighter in the past (at fixed rotational velocity), while the most massive would follow the local relation. However, other groups, such as \\citet{2006MNRAS.366..308B}, find luminosity evolution but no slope change. Moreover, these authors find differences in luminosity evolution ranging from $-$0.2 to $-$2 magnitudes.  To cope with contributions to line widths other than that produced by rotational velocity, \\citet{2006weinerI} measured kinematic line widths ($\\sigma$$_{1D}$), and spatially resolved rotation and dispersion profiles. They combined line--of--sight rotation and dispersion (S$^2_{0.5}$=0.5V$^2_{rot}$+$\\sigma$$^2_{2D}$), and found that this combination correlates well with the integrated line width, demostrating that $\\sigma$$_{1D}$ is a robust velocity indicator, that allows constructing scaling relations with velocity for a population of diverse kinematic properties (dispersion and rotation--dominated galaxies). In another paper, \\citet{2006weinerII} used $\\sigma$$_{1D}$ to study the evolution in the B and J--band TFRs, finding a slope evolution in both bands, and a larger intercept evolution in the B--band ($-$1.5 mag versus $-$0.5 mag in the J--band up to z=1.2). This slope change, which implies a large evolution for the most massive galaxies, is in the opposite sense than that found by \\citet{2002ApJ...564L..69Z} and \\citet{2004A&A...420...97B}. Although the B--band TFR derived by \\citet{2006weinerII} using resolved rotation velocities also supports an intercept evolution, the sample is too small and noisy to measure slope evolution. In addition, \\citet{2007ApJ...660L..35K} demostrated that the scatter in the stellar--mass TFR is lower when using S$_{0.5}$ instead of V$_{rot}$. In \\citet{2009fernandez} we analysed the optical line widths to study the evolution of the TFR in B, R and I--bands, demostrating that all optical emision lines can be used for determining galaxy kinematics. Also, we constructed our local TFRs from data derived in a similar way than the sample of high redshift galaxies, and found evidence of luminosity evolution in all three bands for the greatest redshift range of our sample, 1.1$<$z$<$1.3.  In contrast to the variety of results found in the optical bands, there seemed to be a consensus on the absence of evolution in the K--band TFR \\citep{2005ApJ...628..160C,2006A&A...455..107F}. However, in a recent work, \\citet{2008A&A...484..173P} do find K--band evolution in the sense that galaxies had been fainter in the past, a result opposite to what is found in the optical bands \\citep[for example,][]{2006MNRAS.366..308B}. The disagreement between this work and that of \\citet{2006A&A...455..107F} (both use 3D spectroscopy for deriving velocities) is due to the method for correcting the rotation velocity and, above all, the local relation used as reference. Several works have studied the local TFR in the K--band finding different results, with a slope ranging from $-$6.88 \\citep{2007hammer} to $-$11.3 \\citep{2001ApJ...563..694V}. Then, to obtain reliable conclusions on the evolution of the infrared TFR, is crucial to find the reason of these disagreements, fixing the local relation. Recently, \\citet{2008AJ....135.1738M} presented a universal calibration of the TFR in the 2MASS J, H and K--bands. They showed that in all three bands the relation is steeper for later--type spirals, and they obtained a slope of $-$10.017 in the K--band correcting all galaxies to Sc type.  In this era of precision cosmology, the increasingly large, deep and accurate galaxy surveys within which the DEEP2 project \\citep{2003SPIE.4834..161D,2007ApJ...660L...1D} provides the higher resolution galaxy spectra currently available, allow going one step forward on the study of the TFR: its colour evolution. Models of disc galaxy formation predict different colour evolution depending whether the evolution is mainly due to collapse or accretion \\citep[see, for example,][]{2002A&A...389..761W}. In our previous work \\citep{2009fernandez}, we studied for first time the (R$-$I) TFR evolution up to z$<$1.3, but the dispersion prevented us from distinguishing between one galaxy model from another.  In the present work, we have increased the sample to the largest redshift range 1.1$<$z$<$1.4 for all the Extended Groth Strip (EGS) with DEEP2 spectra available, with the aim to confirm or not the {\\it prima facie} evidence for evolution found previously \\citep{2009fernandez}. Also, we will extend our study to the V \\& K$_{\\rm S}$--bands, in addition to the B, R \\& I--bands, in order to study the most probable galaxy formation model, since the \\citet{2002A&A...389..761W} models predict larger colour differences in these bands. The local sample (0.1$<$z$<$0.3) has been extended to the whole EGS as well. Using this new local sample, we expect to set a more reliable TFR slope at z=0 to compare with our high--redshift sample.  This paper is organised as follows. In Sect. 2, a description of the data and the sample selection criteria are provided. The study of k--corrections and derivation of the absolute magnitudes and rotation velocities are described in Sect. 3. The results are presented in Sect. 4, whereas the last two sections provide the discussion of the results and conclusions. Throughout this article, the concordance cosmology with ${\\Omega}_{\\rm \\Lambda0}=0.7$, ${\\Omega}_{\\rm m0}=0.3$ and $\\rm H_0=70 \\rm \\ km \\rm \\ s^{-1} \\rm \\ Mpc^{-1}$ is assumed. All magnitudes are in the AB zero--point system. ",
        "conclusions": "We have studied the evolution of the TFR in B, V, R, I and K$_{\\rm S}$--bands, using two sample of galaxies in the GSS, one in the redshift range 0.1$<$z$<$0.3 as representative of local galaxies; and the high redshift sample in the redshift range 1.1$<$z$<$1.4. The rotation velocity of all 241 spiral galaxies were measured from optical lines widths, using DEEP2 spectra. Morphology was obtained from HST images, and the absolute magnitudes were derived from a carefully study of the k--correction made with a large set of photometric information. The results of this study can be summarized as follows:  We analized four sets of templates and we concluded that the SED of a spiral galaxy is better reproduced by the nonnegative linear combination of five templates based on the \\citet{2003MNRAS.344.1000B} stellar evolution synthesis codes obtained by {\\tt kcorrect}. However, when the data include a noisy band, the k--corrections calculated by this code are not suitable. Also, we find that the most reliable k--correction is obtained from information in a observed band that roughly match the rest--frame band. In this case, the rest--frame magnitude in the noisy band is better calculated directly from the best--fit template.  When the observed photometry does not match the rest--frame bands, it is possible to calculate the rest--frame magnitudes via an interpolation method, as long as enough information at larger and shorter wavelengths is available. However, an extrapolation method is not reliable for galaxies at z$\\sim$1.25 since the results can vary as more than one magnitude mainly due to a worse fit of the template.  New local TFRs were constructed increasing the sample of our previous work \\citep{2009fernandez} in the redshift range 0.1$<$z$<$0.3. We obtained a similar difference in the zero--point of the TFR for all bands when we compare with the \\citet{2001ApJ...563..694V} local relations, that we attributte to the ratio between rotation velocity and line width.  We derived the high redshift 1.1$<$z$<$1.4 TFRs fixing the slope to the local one in each band, to study the evolution of the zero--points. We confirm the evolution in the B--band TFR found in our previous work \\citep{2009fernandez}, in the sense that galaxies were brighter in the past, and no evolution in the K$_{\\rm S}$--band TFR. Furthermore, for the other bands we find a gradual evolution, with a change in luminosity more noticiable as the band is bluer.  We study the slope of the local K$_{\\rm S}$--band TFR comparing with Millenium simulated data. We obtained that the slope changes depending on the luminosity range used in the fit. Considering this effect, we calculated new local K$_{\\rm S}$--band TFR from the Millenium data in the same luminosity range than our high redshift sample, confirming again no evolution in the K$_{\\rm S}$--band TFR.  Assuming the change in the stellar mass--to--light ratio M/L$_K$ found by \\citet{2007A&A...465..417W}, consistent with that of \\citet{2004ApJ...608..742D}, but no evolution in the K$_{\\rm S}$--band, we obtain that galaxies would have duplicated their stellar mass in the last 8.6 Gyr. Likely, this growth is mainly due to the gas been transformed into stars, as supported by the change in the gaseous O/H abundance found by \\citet{2008A&A...492..371R}. From the mass--to--light ratios derived by \\citet{2007A&A...465..417W} for the optical bands (IMF of Salpeter), and assuming this change in stellar mass, we infer an evolution in the optical luminosity similar to that found in the present work. However, the luminosity and colour evolution simulated by \\citet{2007A&A...465..417W} is not consistent with that found here and would imply a growth of $\\sim$6 in the stellar mass from z=1.25 to z=0.  Therefore, the galaxies, despite having doubled their stellar mass in the last 8.6 Gyr, are nowadays fainter and redder in the optical bands, which can be attributed to an ageing of the stellar populations as consequence of the star formation density decrease observed in the last 8.6 Gyr. This is supported by the change found in the (V--K$_{\\rm S}$) and (R--I) colours (for a fixed velocity), in the sense that galaxies were bluer in the past."
    },
    "1005/1005.3498_arXiv.txt": {
        "abstract": "We present the frequency resolved energy spectra (FRS) of the low-mass X-ray binary dipper \\thirteen\\ during persistent emission in four different frequency bands using an archival \\XMM observation. FRS method helps to probe the inner zones of an accretion disk. \\thirteen\\ is an atoll source  and a type-I burster. We find that the FRS is well described by a single blackbody component with kT in a range 1.0-1.4 keV responsible for the source variability in the frequency ranges 0.002-0.04 Hz, and 0.07-0.3 Hz. We attribute this component to the accretion disk and possibly emission from an existing boundary layer supported by radiation pressure. The appearance of the blackbody component in the lower frequency ranges and disappearance towards the higher frequencies suggests that it may also be a disk-blackbody emission. We detect a different form of FRS for the higher frequency ranges 0.9-6 Hz and 8-30 Hz  which is modeled best with a power-law and a Gaussian emission line at 6.4$^{+0.2}_{-0.3}$ keV with an equivalent width of 1.6$^{+0.4}_{-1.2}$ keV and 1.3$^{+0.7}_{-0.9}$ keV for the two frequency ranges, respectively. This iron fluorescence line detected in the higher frequency ranges of spectra shows the existence of reflection in this system within the inner disk regions. The conventional spectrum of the source also shows a weak broad emission line around 6.6 keV (Boirin et al. 2005). In addition, we find that the 0.9-6 Hz frequency band shows two QPO peaks at 1.4$^{+1.0}_{-0.2}$ Hz and 2.8$^{+0.2}_{-0.2}$  Hz at about 2.8-3.1 $\\sigma$ confidence level. These are consistent with the previously detected $\\sim$ 1 Hz QPO from this source (Jonker et al. 1999). We believe they relate to the reflection phenomenon. The emission from the reflection region, being a variable spectral component in this system, originates from the inner regions of the disk with a maximum size of 4.7$\\times 10^9$ cm and a minimum size of 1.6$\\times 10^8$ cm calculated using light travel time considerations and our frequency resolved spectra. ",
        "introduction": "\\thirteen\\ is an X-ray burster that shows varying intensity dips that repeat with the orbital period. The source is discovered by {\\it Uhuru} and {\\it Ariel V} (Forman et al. 1978; Warwick et al. 1981) and the dips and bursts are recovered with EXOSAT (van der Klis et al. 1985; Parmar et al. 1989). The typical dips last about 30\\% of the orbital cycle, and the 0.3-12.0 keV intensity varies creating sharp deep features in time with about $\\sim$70\\% decrease in the persistent emission count rate. The source is viewed at $i = $ 60--80$\\degmark$ (Frank et al. 1987). The dip recurrence interval is not precisely known, with the best measurement of 2.94(2)~hr (Balucinska-Church et al. 1999). The source is found to have  persistent $\\sim$1~Hz quasi-periodic oscillations (QPOs) (Jonker et al. 1999).  A $BeppoSAX$ observation of \\thirteen\\ between 1.0--150~keV reveals persistent emission of the source  modeled by a cutoff power-law with $\\alpha$ = $1.48 \\pm 0.01$ and \\ecut\\ = $44.1 \\, ^{+5.1} _{-4.4}$~keV together with a blackbody with $kT$ = $1.77 \\pm 0.25$~keV (Balucinska-Church et al. 1999). A more recent $INTEGRAL$ time-averaged 4.0-200.0~keV spectrum (combined JEM-X and ISGRI spectra) of \\thirteen\\ can be best fitted using a blackbody emission of $kT$ = $1.7\\pm 0.3$~keV and a Comptonized plasma emission model (CompTT, Titarchuck 1994) of $kT$ = $196.4\\pm 28.5$ keV with a seed in photon temperature of $0.4\\pm 0.2$~keV, and a Compton optical depth to scattering of $\\tau$ = 0.002-0.007 (Balman  2009). An $RXTE$ observation below 20 keV reveals a cutoff power-law with $\\alpha$ = $1.23 \\pm 0.07$ and a blackbody with $kT$ = $1.36 \\pm 0.06$~keV for the persistent emission (Barnard et al. 2001). A recent $Suzaku$ observation between 0.8 and 70 keV has a non-dip spectrum consistent with a blackbody model of $kT$ = $1.35 \\pm 0.36$~keV and a power-law with an $\\alpha$ = $1.67 \\pm 0.1$ (Balucinska-Church et al. 2009). Boirin et al. (2005) detect Fe XXV and Fe XXVI absorption features from this source using the same archival observation analyzed in this paper. They find changes in the properties of the Fe XXV and Fe XXVI absorption features from persistent to dipping intervals indicating the presence of less-ionized material in the line of sight during dips (ionization parameter $\\xi$ decreases from log($\\xi$) of $3.9\\pm 0.1$ to log($\\xi$) of $3.13\\pm 0.07$).  It is suggested that absorption lines in dipping low-mass X-ray binaries are a result of a circularly symmetric photoionized region on the outer disk and as the source is viewed from the absorber at all times, the variations in the ionization parameter of this region causes the different levels of dipping in these systems (see also Diaz-Trigo et al. 2006, 2009). Church et al. (2005) has studied the same data set as Boirin et al. (2005) using absorption line ratios of Fe XXV/Fe XXVI and curve of growth analysis leading to the conclusion that a collisionally ionized plasma with kT=31 keV can reproduce the absorption  lines in the \\XMM Spectra. As they find that this temperature is close to the accretion disk corona (ADC) temperature (they have calculated using the $BeppoSax$ spectrum), they propose that the absorption lines in dipping low-mass X-ray binaries are produced in the ADC and there is no need to invoke a separate location where a highly ionized absorber exists.  In this paper, we aim at deriving frequency resolved spectra of \\thirteen\\ in an attempt to reveal different emission regions and  absorbing regions in the system extracting the variable component of the emission spectrum as opposed to the non-variable component by the described method in the paper. We expect that this method should, also, reveal the location of the emitting and absorbing regions in similar (dipping) LMXBs. ",
        "conclusions": "We have analyzed the frequency resolved energy spectra (FRS) of the low-mass X-ray binary dipper \\thirteen\\ during persistent emission in four different frequency bands using an archival \\XMM observation. We recover two different continuum shapes for the four frequency ranges we have studied. We find that the FRS shows a single blackbody component with kT in a range 1.0-1.4 keV for the variability in the frequency ranges 0.002-0.04 Hz, and 0.07-0.3 Hz. We attribute this component to the accretion disk and possibly emission from an existing boundary layer supported by radiation pressure. On the other hand, we stress that the emergence of a blackbody component in the lower frequency ranges and disappearance towards the higher frequencies suggests that it may also be a disk-blackbody emission given the FRS properties derived for Z-sources. We find a different form of FRS for the frequency ranges 0.9-6 Hz and 8-30 Hz  which is a power-law model together with a Gaussian emission line at 6.4$^{+0.2}_{-0.3}$ keV with an equivalent width of 1.6$^{+0.4}_{-1.2}$ keV and 1.3$^{+0.7}_{-0.9}$ keV for the two frequency ranges, respectively. The width $\\sigma$ of the line is $<$ 0.5 keV and 0.7$^{+0.3}_{-0.5}$ for the 0.9-6 Hz and 8-30 Hz ranges. In addition, we recover different neutral hydrogen column density for the FRS derived in our work. The $N_{\\rm H}$ is highest in the 0.9-6.0 Hz range with a value of 2.8$^{+0.6}_{-0.8}$ $\\times 10^{22}$ cm$^{-2}$ and is only slightly lower for the 8.0-30.0 Hz range, but lower by a factor of 3-4 in the 0.002-0.04 Hz and 0.07-0.3 Hz ranges. So, there is more associated absorption towards the inner disk and a distinct absorption difference for the two different continuum detected for the FRS. We find that the emission structure within the inner regions of the disk is such that there is either boundary layer emission and an embedded reflection region with higher absorption towards the inner regions or there is a reflection region and an ADC in the inner disk followed by a disk-blackbody towards the outer disk regions. The derived iron fluorescence line shows the existence of reflection/reprocessing in this system within the inner disk region. The size of the reflection region is in a range 4.7-0.16$\\times 10^9$ cm assuming light travel times. We suggest that the FRS method applied in this work have recovered the reflected component in the system that had not been found before by conventional spectral methods. We stress that FRS can be used in this manner to find reflection components in other LMXBs where conventional spectra can not be adequate. Finally, we find that the 0.9-6 Hz frequency band shows two QPO peaks at 1.4$^{+1.0}_{-0.2}$ Hz and 2.8$^{+0.2}_{-0.2}$  Hz at about 2.8-3.1 $\\sigma$ confidence level with integrated rms variability of 15$\\%$ and 11$\\%$, respectively. We  relate the origin of the QPOs to the reflection phenomenon. The width of the 6.4 keV line increases with increasing frequency where the location, the QPOs are produced, is the outer most regions of the reflection zone with the least spread in the line and the highest $N_{\\rm H}$, thus, possibly a cold reflection zone (i.e., region with large clumps)."
    },
    "1005/1005.3667_arXiv.txt": {
        "abstract": "Observations of the polar region of the Sun are critically important for understanding the solar dynamo and the acceleration of solar wind. We carried out  precise magnetic observations on both the North polar region and the quiet Sun at the East limb with the Spectro-Polarimeter of the Solar Optical Telescope aboard {\\it Hinode} to characterize the polar region with respect to the quiet Sun. The average area and the total magnetic flux of the kG magnetic concentrations in the polar region appear to be larger than those of the quiet Sun. The magnetic field vectors classified as {\\it vertical} in the quiet Sun have symmetric histograms around zero in the strengths, showing balanced positive and negative flux, while the histogram in the North polar region is clearly asymmetric, showing a predominance of the negative polarity. The total magnetic flux of the polar region is larger than that of the quiet Sun. In contrast, the histogram of the {\\it horizontal} magnetic fields is exactly the same between the polar region and the quiet Sun. This is consistent with the idea that a local dynamo process is responsible for the horizontal magnetic fields. A high-resolution potential field extrapolation shows that the majority of magnetic field lines from the kG-patches in the polar region are open with a fanning-out structure very low in the atmosphere, while in the quiet Sun, almost all the field lines are closed. ",
        "introduction": "The polar regions of the Sun have some distinct properties compared with those of the quiet Sun: (1) the coronal holes in the polar regions are extended and stationary, (2) the fast solar wind emanates from the polar region, (3) the meridional flow may reach the polar regions, which must have a sink for the flow coming to the polar regions, and (4) the global poloidal magnetic field of the Sun manifests itself in the polar regions. Therefore, information on the magnetic properties of the polar regions is critically important for the understanding of the solar dynamo problem. The acceleration mechanism of the fast solar wind must also be related to the spatial and temporal structures of magnetic fields on the surface of the Sun, and the physical process creating the wind is still a topic of debate  \\citep{Wang89, Wang90, McCo00, Kojima01, Tu05}.  The polar regions have been observed with ground-based magnetograms \\citep{Babcock55, Severny71, svalgaard78, Tang91, Lin94, Homann97, Fox98, Oku04, Blanco07}. These observations as well as observations with {\\it SOHO/MDI} \\citep{Bene04} have provided measurements of the line-of-sight magnetic component. For more accurate observations, full Stokes polarimetry for the polar regions has been carried out \\citep{Lites96}, but because of the strong intensity gradient and the foreshortening effect at the solar limb, the spatial resolution of the ground-based measurements has been compromised. The magnetic properties of the polar region are poorly understood.  \\citet{tsuneta08} found with {\\it Hinode} \\citep{Kosugi07} that there are patchy concentrations of magnetic fields with field strengths exceeding 1 kG in the polar region, confirming the result of \\citet{Oku04}. These magnetic fields are almost {\\it vertical} to the local surface of the Sun. Ubiquitous magnetic concentrations with field strengths larger than 1 kG are known to exist in the quiet Sun as well \\citep[e.g.,][]{david07a}. This paper focuses on determining whether there is any difference in the magnetic properties between the polar region and the quiet Sun. If there is a difference, what is the implication to the problems of solar dynamo and the acceleration of the fast solar wind? With this motivation, we observed the polar region and the quiet Sun as a fiducial data for the analysis of the polar region with the spectropolarimeter (SP) \\citep{Lites01} of the Solar Optical Telescope (SOT) \\citep{tsuneta08a, Suematsu08, Ichi08, Shimizu08} aboard {\\it Hinode}.  The sensitivity to a given magnetic field vector depends on its orientation with respect to the observer and on the Stokes parameter used to detect it. Stokes {\\it Q} and {\\it U} signals are used to detect the field component perpendicular to the line-of-sight, while Stokes {\\it V} signal is used for the field component parallel to the line-of-sight. Measurement through Stokes {\\it V} is more sensitive than that through Stokes {\\it Q} and {\\it U} when the field is weak. We observed the quiet Sun located at the East limb instead of the quiet Sun at the disk center, so that the sensitivity to detect any weak magnetic field is the same for both the polar region and the quiet Sun. This allows a direct comparison between the polarimetric data for the polar region with that for the quiet Sun.  \\begin{figure} \\begin{center} \\epsscale{.80} \\plotone{f1.eps} \\end{center} \\caption{ Magnetic landscape of the polar region of the Sun; The map of the magnetic field strength on the sky plane is converted to the map seen from above the North pole. The observations took place at 00:10---07:26 UT on 2007 September 25. East is to the left. The pixel size is $0.16''$. The magnetic field strengths were obtained only for pixels whose polarization signal exceeds a given threshold (see text). The field of view is $320''$ by $477.6''$. The size of the field of view for North-South direction ($163.84''$) is expanded to $477.6''$ as a result of correction for foreshortening. The circular arc is the latitudinal line per $5^\\circ$ from the North pole. } \\label{npole2} \\end{figure}  \\begin{figure} \\begin{center} \\epsscale{1.} \\plotone{f2.eps} \\end{center} \\caption{ Two solutions of the magnetic field vectors $\\sharp$1 and $\\sharp$2 with the inclination angles of ${\\it i_{\\rm 1}}$ and ${\\it i_{\\rm 2}}$ with respect to the photospheric normal for the two possible azimuth angles of $\\chi$ and $\\chi + \\pi$ due to the $180^\\circ$ ambiguity, respectively. $\\theta$ is the angle between the local normal and the line-of-sight (L.O.S.), $\\gamma$ the inclination ($ 0 \\leq \\gamma \\leq \\pi $) of the magnetic field vector with respect to L.O.S., and $\\chi$ is the azimuth angle ($ 0 \\leq \\chi \\leq \\pi $). Panel $(a)$ shows these angles on the sky plane. Note that the lines indicating angles are represented by the straight lines, not by the segments of arc, for simplicity. Panel $(b)$ indicates the same information in terms of vectors at the position of any observing pixel. The dotted and dashed lines represent Great-circles connecting the L.O.S. direction and the photospheric normal and the magnetic vector $\\sharp$1, respectively. Panel $(c)$ is the same as the Panel $(b)$ but for the magnetic vector $\\sharp$2. These two representations shown in the panel $(a)$, and the panels $(b)$ and $(c)$ are equivalent. } \\label{incli1} \\end{figure}  \\begin{figure} \\begin{center} \\epsscale{1.0} \\plotone{f3.eps} \\end{center} \\caption{ Scatter plot of the two solutions ${\\it i_{\\rm 1}}$ and ${\\it i_{\\rm 2}}$ of the inclination angles with respect to the local normal to the photosphere for the North pole. The color table indicates the density of the data points per $1^\\circ \\times 1^\\circ $. Magnetic fields are classified as {\\it vertical} to the local surface if they are located either close $0^\\circ$ or $180^\\circ$ (dashed lines), and as {\\it horizontal} if located near $90^\\circ$ (bold lines). For the concentration along the line ${\\it i_{\\rm 1}} \\sim {\\it i_{\\rm 2}}$ as indicated by the white line, see text. } \\label{incli2} \\end{figure}  \\begin{figure} \\begin{center} \\epsscale{.80} \\plotone{f4a.eps} \\end{center} \\caption{ {\\it (a)} Map of signed strength of the magnetic field vectors classified as {\\it vertical} in the polar region. The original data is the same as that of Fig. \\ref{npole2}. The pixel size is $0.16''$. The magnetic field strength was obtained only for pixels whose polarization signal exceeds a given threshold (see text). These panels can be directly compared with the quiet Sun maps (Fig. \\ref{elimb3}). The scale size and the color table for these maps are the same. } \\label{npole3} \\end{figure} \\begin{figure} \\begin{center} \\epsscale{.80} \\plotone{f4b.eps} \\end{center} Fig. 4 ---{\\it (b)} same as Figure 4 {\\it (a)} but map of field strength of magnetic field vector classified as {\\it horizontal}. \\label{npole3a} \\end{figure}  \\begin{figure} \\begin{center} \\epsscale{1.00} \\plotone{f5.eps} \\end{center} \\caption{ Magnetic landscape of the quiet Sun near the East limb; The map of the magnetic field strength on the sky plane is converted to the map seen from above the East limb. The observations are done at 21:00---23:16 UT on 2007 November 28. North is up. The pixel size is $0.16''$. The field of view is $100''$ (East-West) by $162.84''$ (North-South). The size of the field of view for East-West direction ($100''$) is expanded to  $397.6''$ as a result of correction for foreshortening. The circular arc is the latitudinal line per $5^\\circ$ from the East limb. The magnetic field strengths were obtained only for pixels whose polarization signal exceeds a given threshold (see text). This figure can be directly compared with the polar map (Fig. \\ref{npole2}). The scale size and the color table for these maps are the same. } \\label{elimb2} \\end{figure}  \\begin{figure} \\begin{center} \\epsscale{1.00} \\plotone{f6.eps} \\end{center} \\caption{ {\\it (a)} Map of signed strength of the magnetic field vectors classified as {\\it vertical} near the East limb. {\\it (b)} Map of strength of the magnetic field vectors classified as {\\it horizontal} near the East limb. The original data is the same as that of Fig. \\ref{elimb2}. The pixel size is $0.16''$. The magnetic field strengths were obtained only for pixels whose polarization signal exceeds a given threshold (see text). These panels can be directly compared with the polar maps (Fig. \\ref{npole3}). The scale size and the color table for these maps are the same. } \\label{elimb3} \\end{figure}  \\begin{figure} \\begin{center} \\epsscale{0.9} \\plotone{f7.eps} \\end{center} \\caption{ Data from the kG magnetic concentrations arbitrarily chosen in the North polar region (Fig. \\ref{npole2}); the histograms of intrinsic magnetic fields strength (left), maps of intrinsic magnetic field strength (center) and the filing factor (right) seen from above the pole. The pixel size of the maps is 0.16$''$. } \\label{npatch} \\end{figure}  \\begin{figure} \\begin{center} \\epsscale{0.9} \\plotone{f8.eps} \\end{center} \\caption{ Data from the  kG magnetic concentrations arbitrarily chosen in the quiet Sun (Fig. \\ref{elimb2}); the histograms of intrinsic magnetic fields strength (left), maps of intrinsic magnetic field strength (center) and the filing factor (right) seen from above the East limb. The quiet Sun has the kG-patches with both positive and negative magnetic polarities, and the absolute intrinsic magnetic field strengths are shown. The pixel size of the maps is 0.16$''$. } \\label{epatch} \\end{figure} ",
        "conclusions": "In this paper, we have made a detailed comparison of the photospheric magnetic properties between the polar region and the quiet region. We found that the magnetic properties of the polar region are significantly different from those of the quiet Sun. Positive and negative magnetic fields are well balanced in the quiet Sun, while in the North polar region, negative-polarity magnetic fields dominate the other polarity. The excess negative magnetic field makes both total unsigned magnetic flux and net (signed) magnetic flux larger than those of the quiet Sun.  \\subsection{Kilo-Gauss Magnetic Patches} All the large magnetic concentrations in the polar region essentially have the same magnetic polarity, while the smaller patches have the same and the opposite magnetic polarities. \\citet{Shimojo09} reported bipolar emergence of magnetic fields in the polar region, and this may contribute to the conspicuous minority polarity patches. The average area and the average magnetic flux of the kG patches appear to be considerably different between the two regions: the sample data set tabulated in Table 1 shows that kG magnetic concentrations in the polar region have a factor of 3.3 larger average area, and have a factor of 4.8 larger total magnetic flux than those in the quiet Sun. The number of the kG-patches in the polar region appears to be larger than that of the quiet Sun.  The origin and/or evolution of the kG-patches would be different between the polar region and the quiet Sun. The different properties of the kG-patches in the polar regions and the quiet Sun could be due to the different environment in which they evolve. There may be a higher chance in the quiet Sun that the positive and negative patches collide, reconnect, and lose magnetic energy or submerge as a result. As such, the environment in the quiet Sun may not allow the elemental magnetic concentrations to grow.  \\subsection{Horizontal magnetic fields} From Fig. \\ref{PDF}, 98\\% of the horizontal magnetic fields has the intrinsic field strengths smaller than 700 G both in the quiet Sun at the limb and the polar region. Here the horizontal magnetic fields refer to the magnetic field vectors classified as horizontal. \\cite{ishi09} reported that 93\\% of horizontal magnetic fields have intrinsic field strengths smaller than 700 G, and 98\\% smaller than 1 kG in the quiet Sun and a plage region both located near the center of the solar disk. Magnetic field strength of 700 G corresponds to the typical equipartition field strength at the level of granules, where the density is $10^{-6}$ g cm$^{-3}$ at the depth of 500 km and the velocity is 2 km s$^{-2}$. If we take the mean density $2.7 \\times 10^{-7}$ g cm$^{-3}$ at the base of the photosphere, the equipartition field strength is 350 G. Even in this case, 92\\% of the horizontal magnetic fields has the intrinsic field strengths smaller than 350 G both in the quiet Sun at the limb and the polar region. Although the vertical magnetic flux in the polar region is larger than that of the quiet Sun, the horizontal magnetic flux in the polar region is remarkably similar to that of the quiet region. The small difference between the two results (98\\% in the present paper and 93\\% in \\cite{ishi09} at 700 G) could be due to the following two factors: (1) A higher atmosphere is observed in the limb observations, so that the measured magnetic field strength may be smaller in the limb observations. (2) The horizontal magnetic fields are observed through Stokes Q and U in the disk center observations, while Stokes V is involved in the limb observations. Thus, the limb observations may have higher sensitivity to the horizontal magnetic fields.  \\cite{ishi09} proposed that a local dynamo process be uniformly operating in the photosphere or below based on the observations that the probability distribution function of the transient horizontal magnetic fields in the quiet Sun is the same as that in a plage region. If we assume that the magnetic properties of the quiet Sun in the disk center analyzed by \\cite{ishi09} are the same as those of the quiet Sun at the East limb analyzed in this paper, we would conclude that the distribution function of the horizontal magnetic fields and the amount of the magnetic flux are exactly the same in the quiet Sun, a plage region and in the polar region. This provides firmer evidence that a local dynamo process driven by the granular motion is responsible for the generation of the horizontal magnetic fields.  \\subsection{Polar coronal holes} \\citet{Kano08b} obtained the temperatures and the densities of the polar coronal hole and the quiet Sun with the X-ray telescope (XRT) \\citep{Golub07,Kano08a} aboard {\\it Hinode}. These are the most accurate measurements carried out using  the broadband filters by correcting for the effect of scattered X-rays with the lunar occultation. The temperatures and electron densities then derived are 1.0 MK and 1.0$-$1.5 $ \\times 10^8 {\\rm cm^{-3}}$ for the polar coronal hole, and 1.5$-$2.0 MK and 2.0$-$3.0 $\\times 10^8 {\\rm cm^{-3}}$ for the quiet Sun. The plasma pressures are 1.4$-$2.1 $\\times 10^{-2} {\\rm dyne \\cdot cm^{-2}}$ and 4.1$-$8.3 $\\times 10^{-2} {\\rm dyne \\cdot cm^{-2}}$ for the coronal hole and the quiet Sun, respectively. Thus, plasma pressure in the coronal hole is a factor of 3$-$6 smaller than that of the quiet Sun.  The coronal hole in the upper corona inside the Alfv\\'enic sphere should have collapsed due to the lateral pressure from the surrounding quiet Sun corona, if the magnetic field strength in the polar coronal hole is much smaller than that of the quiet corona. The polar coronal hole is stable for a long time, and the coronal hole and the surrounding corona of the quiet Sun must be in pressure equilibrium. In this paper, we showed that the magnetic flux of the polar region in the photosphere is larger than that in the quiet Sun (Fig. \\ref{PDF}). Indeed, magnetic pressure estimated from Fig. \\ref{BrBh} is 1 dyn cm$^{-2}$ in the polar region and $3 \\times 10^{-3}$ dyn cm$^{-2}$ above the quiet Sun at the height of $7 \\times10^4$ km. The boundary of the polar coronal hole and the quiet Sun must be determined by the total pressure balance between the coronal hole and the corona in the quiet Sun, and the polar corona expands to the lower latitude until it reaches the magnetic plus plasma pressure equilibrium with the quiet Sun corona. This may be the reason why the fast solar wind from the polar region reaches the lower latitude \\citep{McCo00}.  \\subsection{Polar coronal activities} \\citet{Defor97} investigated the correlation between solar plumes and polar magnetic field with {\\it SOHO}, and found that there is opposite magnetic flux close to the foot-point of solar plumes. In the polar coronal hole, X-ray jets are observed with high occurrence rate (60 polar X-ray jets per day on average) \\citep{Kamio07, Cirtain07,Sav07}. \\citet{Shimojo09} revealed that the opposite magnetic fields close to the kG-patches are related to the occurrence of the X-ray jets. Distribution of such opposite magnetic fields around a kG-patch is indeed seen in the (signed) magnetic field map in the North polar region (Fig. \\ref{npole3}), and may be responsible for various coronal activities."
    },
    "1005/1005.3673_arXiv.txt": {
        "abstract": "The stability properties of the Einstein Static solution of General Relativity are altered when corrective terms arising from modification of the underlying gravitational theory appear in the cosmological equations. In this paper the existence and stability of static solutions are considered in the framework of two recently proposed quantum gravity models. The previously known analysis of the Einstein Static solutions in the semiclassical regime of Loop Quantum Cosmology with modifications to the gravitational sector is extended to open cosmological models where a static neutrally stable solution is found. A similar analysis is also performed in the framework of Ho\\v{r}ava-Lifshitz gravity under detailed balance and projectability conditions. In the case of open cosmological models the two solutions found can be either unstable or neutrally stable according to the admitted values of the parameters. ",
        "introduction": "The Einstein Static (ES) Universe is an exact solution of Einstein's equations describing a closed Friedmann-Robertson-Walker model sourced by a perfect fluid and a cosmological constant (see, for example, \\cite{Hawking:1973uf}). This solution is unstable to homogeneous perturbations as shown by Eddington \\cite{Edd:1930}, furthermore it is always neutrally stable against small inhomogeneous vector and tensor perturbations and neutrally stable against adiabatic scalar density inhomogeneities with high enough sound speed \\cite{Barrow:2003ni}.  In recent years there has been renewed interest in the ES Universe because of its relevance for the Emergent Universe scenario \\cite{Ellis:2002we} in which the ES solution plays a crucial role, being an initial state for a past-eternal inflationary cosmological model. In the Emergent Universe scenario the horizon problem is solved before inflation begins, there is no singularity, no exotic physics is involved, and the quantum gravity regime can even be avoided. This model, relying on the choice of a particular initial state, suffers from a fine-tuning problem which is ameliorated when modifications to the cosmological equations arise but then a mechanism is needed to trigger the expanding phase of the Universe (see \\cite{Lidsey:2006md}).  The existence of ES solutions along with their stability properties has been widely investigated in the framework of General Relativity for several kinds of matter \ufb01elds sources (\\cite{Barrow:2009sj} and references therein). ES solutions also exist in several modified gravity models \\cite{Boehmer:2010xa} ranging from the Randall-Sundrum braneworld scenario \\cite{Gergely:2001tn} to Gauss-Bonnet modified gravity and $f(R)$ theories \\cite{Clifton:2005at}. The issue of the existence and stability of ES solutions has also been considered in the semiclassical regime of Loop Quantum Cosmology (LQC), in either the case of correction to the matter sector \\cite{Mulryne:2005ef} or the case of correction to the gravitational sector \\cite{Parisi:2007kv}. Recently the same issue has been also considered in the framework of Ho\\v{r}ava-Lifshitz (HL) gravity \\cite{Wu:2009ah} and IR modified Ho\\v{r}ava gravity \\cite{Boehmer:2009yz,Park:2009zra}.  When dealing with higher order modified cosmological equations, the existence of many new ES solutions is possible, whose stability properties, depending on the details of the single theory or family of theories taken into account, are substantially modified with respect to the classical ES solution of General Relativity (GR).  Often in such analysis the case of closed $(k=1)$ cosmological models is the only one considered, neglecting the intriguing possibility of static solutions in open $(k=-1)$ cosmological models. Here we point out that, due to the aforementioned corrections to the cosmological equations, open ES models may be found even in the case of a vanishing cosmological constant or when the perfect fluid has vanishing energy density. Examples are explicitly provided for the case of LQC and HL gravity.  This paper is structured as follows. In Sec. II, we consider static solutions in the framework of LQC, following and enlarging the analysis already performed in \\cite{Parisi:2007kv}. It is shown that, beside the ES solution of GR, a LQC solution arises also in the case of open cosmological models; its stability is completely characterized. Following the same approach, in Sec. III we consider static cosmological solutions in the context of HL gravity with detailed balance and projectability condition. Two solutions are found along with their stability properties. In Sec. IV, some conclusions are eventually drawn. ",
        "conclusions": "We have considered the existence of static solutions in the framework of two recently proposed quantum gravity models, namely, LQC and HL gravity. We have shown that the inclusion of a negative curvature index $k=-1$ enlarges the ranges of existence of the solutions affecting their stability properties thus providing new interesting results. The solutions found display stability conditions rather different from those of the corresponding solutions in closed models and from the stability properties of the standard ES solution of GR.  In the case of LQC gravitational modifications to the Friedmann equations, a negative curvature index allows a neutrally stable static solution with $\\Lambda<\\kappa\\rho_{c}$ and $w<-1$, in contrast to the GR case. In particular the LQC static solution exists and is stable in the case of positive values of the cosmological constant as long as  $\\Lambda< \\kappa \\rho_{c}$.  In the case of HL gravity two static solutions are found. The inclusion of the negative curvature index leads to a static solution ($HL1$) with negative cosmological constant and vanishing energy density which is neutrally stable against homogeneous perturbations. Furthermore, a negative curvature index allows a static solution ($HL2$) which can be either a saddle, for $w<-1$ and $w>1/3$, or a center for $-1<w<1/3$.  As already observed in the frameworks of different modified models \\cite{Lidsey:2006md, Mulryne:2005ef, Parisi:2007kv}, the regime of infinite cycles about the center fixed points must be eventually broken in order to enter the current expanding universe phase. To this aim a further mechanism is needed, whose analysis is beyond the scope of this paper."
    },
    "1005/1005.4053_arXiv.txt": {
        "abstract": "We estimate large-scale curvature perturbations from isocurvature fluctuations in the waterfall field during hybrid inflation, in addition to the usual inflaton field perturbations. The tachyonic instability at the end of inflation leads to an explosive growth of super-Hubble scale perturbations, but they retain the steep blue spectrum characteristic of vacuum fluctuations in a massive field during inflation. The power spectrum thus peaks around the Hubble-horizon scale at the end of inflation. We extend the usual $\\delta N$ formalism to include the essential role of these small fluctuations when estimating the large-scale curvature perturbation. The resulting curvature perturbation due to fluctuations in the waterfall field is second-order and the spectrum is expected to be of order $10^{-54}$ on cosmological scales. ",
        "introduction": "Cosmological inflation driven by a single scalar field has many attractively simple properties. In particular the quantum fluctuations of a slow-rolling scalar field have an almost scale invariant power spectrum and these give rise to adiabatic curvature perturbations on super-Hubble scales which have an effectively Gaussian distribution for canonical scalar fields. This is consistent with current observations of primordial curvature perturbations, but there is much interest in possible deviations from this simple paradigm which would require either more than one relevant field during inflation or a non-canonical field, or both. Multiple fields can support non-adiabatic perturbations on super-Hubble scales which can affect the resulting primordial power spectrum and lead to non-Gaussianity.  Hybrid inflation, driven by the energy density of a false vacuum state which is destabilised when a slow-rolling field reaches a critical value, was originally proposed by Linde \\cite{Linde:1991km} and subsequently analysed by Linde \\cite{Linde:1993cn} and many others \\cite{Mollerach:1993sy,Copeland:1994vg}. It has proved to be particularly successful for realising inflation in supersymmetric models of particle physics \\cite{Copeland:1994vg,Dvali:1994ms,Lyth:1998xn,Mazumdar:2010sa}. In hybrid models there is another field, usually called the waterfall field, which is trapped in a false vacuum state until the instability is triggered. Nonetheless the vacuum fluctuations of this field are usually neglected since the field is massive when scales corresponding to large-scale structure of our observable universe leave the horizon. Quantum fluctuations of a massive scalar field, with mass much larger than the Hubble scale, remain overdamped even on super-Hubble scales, corresponding to decaying, oscillating mode functions and thus a steep blue power spectrum.  However it was recently suggested \\cite{Mulryne:2009ci,Levasseur:2010rk} that non-adiabatic large-scale perturbations in the waterfall field could play an important role, in particular leading to non-Gaussian curvature perturbations \\cite{Mulryne:2009ci}, since these long wavelength modes experience the most rapid growth due to the tachyonic instability in the waterfall field at the end of inflation. It is known that a slow transition during hybrid inflation, allowing inflation to continue for some period after the tachyonic instability, could lead to large curvature perturbations on scales which leave the horizon around the time of the instability \\cite{Randall:1995dj,GarciaBellido:1996qt}. But it is difficult to model the large-scale primordial curvature perturbations through the phase transition where perturbations about the classical background necessarily become large.  We reconsider this issue in this paper, using a $\\delta N$-formalism~\\cite{Sasaki:1995aw,Sasaki:1998ug,Wands:2000dp,Lyth:2004gb} to evaluate the primordial curvature perturbation on large scales. We stress the essential role of small, Hubble-scale field perturbations at the end of inflation in determining the local integrated expansion $N=\\int H\\, dt$ in parts of the universe with different values of the waterfall field averaged on large, super-Hubble scales. We show that the variance of the waterfall field becomes dominated by Hubble-scale perturbations when the tachyonic instability begins, rapidly leading to the end of inflation. The duration of inflation is shown to be independent of the large-scale field at first-order, simply due to the symmetry of the potential. The curvature perturbation due to long-wavelength modes of the waterfall field are shown to be suppressed due to the steep blue spectrum of the waterfall field fluctuations, similar to the case of false vacuum inflation supported by thermal corrections~\\cite{Gong:2008ni}. ",
        "conclusions": "We have estimated the primordial curvature perturbation produced by fluctuations in the waterfall field during hybrid inflation. We have calculated linear perturbations about the classical background trajectory, $\\chi=0$, during slow-roll inflation and then studied how these affect the primordial curvature perturbation when the waterfall field is released and inflation comes to an end. To do this we have used a extension of the usual $\\delta N$-formalism that identifies the primordial curvature perturbation with the perturbation in the local expansion on a uniform-density hypersurface, $\\zeta=\\delta N$.  The challenge for the standard $\\delta N$ approach is that the homogeneous classical background solution, with smooth fields on scales much larger than the Hubble-horizon, fails to provide a good description of the dynamics. A classical solution which starts precisely in the false vacuum state stays there and inflation never ends. Thus even a tiny perturbation away from this background solution results in a huge apparent change in the local expansion, $N$. In fact the false vacuum is destabilised mainly due to fluctuations of the waterfall field on scales close to the Hubble-horizon at the end of inflation.  The coupling between the slow-rolling inflaton and the waterfall field that leads to the tachyonic instability also makes the waterfall field massive before the transition. Thus quantum fluctuations of the waterfall field are not amplified during inflation and remain in a quantum state even on super-Hubble scales until the critical value of the inflaton is reached. Once the tachyonic instability is triggered, there is an explosive growth of long-wavelength modes, but they retain the steep blue spectrum on super-Hubble scales. On the other hand much smaller scales, well inside the Hubble horizon at the transition are stabilised by spatial gradients and remain in their vacuum state. The resulting power spectrum for the waterfall field thus peaks on scales around the Hubble scale at the transition, as shown in Figure~\\ref{pu2}, and these modes play an essential role in the dynamics that must be included when calculating the effect of large scale fluctuations in the waterfall field.  Thus we identify the primordial curvature perturbation due to perturbations in the waterfall field, $\\delta\\chi_L$ on some large scale $L$, with the perturbation in the average expansion, $\\langle N_f \\rangle$, {\\em including} fluctuations in the waterfall field on Hubble scales at the transition. We adopt a Gaussian distribution for the field whose average value in a region of size $L$ is $\\delta\\chi_L$, but whose variance is given by $\\sigma_*^2=\\P_u(k_*)/a^2$. We find two main results: \\begin{itemize} \\item from the symmetry of the potential (\\ref{pot}) under $\\chi\\to-\\chi$ we see that the primordial curvature perturbation is independent of $\\delta\\chi_L$ at first order, and the curvature perturbation is second-order in $\\delta\\chi_L$. \\item the spectrum of the primordial curvature perturbation due to fluctuations in the waterfall field on large scales is suppressed by a factor of order $(k_*L)^3$ which for cosmological scales is likely to be of order $10^{-54}$. \\end{itemize} While we have considered the specific hybrid potential (\\ref{pot}) and presented numerical solutions for specific parameter choices we believe the general conclusions will hold for all hybrid models in which the waterfall field is massive during slow-roll inflation and for which the end of inflation occurs due to a rapid tachyonic instability.  It is not surprising that the effect of small scale fluctuations must be taken into account when estimating the primordial curvature perturbation when the large-scale field is very close to zero. This is the case, for example, in the curvaton model where the curvature perturbation, $\\zeta \\sim \\delta\\sigma/\\sigma$, could become very large when the background field, $\\sigma$, is close to zero in some regions of the universe \\cite{Linde:2005yw}. The apparent singular behaviour of the curvature perturbation in this case is regularised by smaller scale fluctuations in the curvaton field~\\cite{Sasaki:2006kq}.  In our analysis we have considered only the early stages of the tachyonic instability. Eventually inflation ends and the coupled fields oscillate about the true vacuum $\\phi=0$ and $\\chi^2=2M^2/\\sqrt\\lambda$. Numerical simulations are required to study resonant particle production and other non-perturbative effects in this regime which goes beyond the scope of this paper. Resonance is often most efficient in long-wavelength modes, but fields which are massive throughout slow-roll inflation necessarily give rise to a steep blue spectrum of perturbations and hence the large-scale power is suppressed with respect to smaller scales \\cite{Liddle:1999hq,Barnaby:2006cq,Barnaby:2006km}.  \\bigskip  {\\em Note added:} While this work was in progress, two papers appeared \\cite{Abolhasani:2010kr,Lyth:2010ch} which study the generation of large scale curvature perturbations at the end of hybrid inflation from complementary viewpoints and whose conclusions appear to be consistent with our results."
    },
    "1005/1005.1730_arXiv.txt": {
        "abstract": "{ We present the first \\hersc\\ PACS and SPIRE images of the low-metallicity galaxy NGC6822 observed from 70 to 500 \\mic\\ and clearly resolve the H~{\\sc ii} regions {\\revised with PACS and SPIRE}. We find that the ratio 250/500 is dependent on the 24 \\mic\\ surface brightness in NGC6822, which would locally link the heating processes of the coldest phases {\\revised of dust in the ISM to the star formation activity}. We {\\revised model the SEDs} of some regions H~{\\sc ii} regions and less active regions across the galaxy and find that the SEDs of H~{\\sc ii} regions {\\revised show} warmer ranges of dust temperatures. We derive {\\revised very} high dust masses when graphite is used in our model to describe carbon dust. Using amorphous carbon, instead, requires less dust mass to account for submm emission due to its lower emissivity properties. This indicates that SED models including \\hersc\\ constraints {\\revised may} require different dust properties than commonly used. The global G/D of NGC6822 is finally estimated to be 186, using amorphous carbon.} ",
        "introduction": "The absorption of stellar radiation and its reemission by dust at infrared (IR) wavelengths is a fundamental process controlling the heating and cooling of the interstellar medium (ISM). The \\iras, {\\it ISO} and \\spitz\\ IR space telescopes launched the studies of the physics and chemistry of dust and gas, revealing their roles in the matter cycle and thermal balance in galaxies. The mid-infrared (MIR) to far-infrared (FIR) wavelength windows provide the necessary observational constraints on the spectral energy distribution (SED) modelling of galaxies from which properties of the polycyclic aromatic hydrocarbons (PAHs) and the warm ($>$30K) and hot dust can be determined. Now, the \\hersc\\ Space Observatory \\citep{Pilbratt2010} is probing the submillimeter regime, where the coldest phases ($<$30K) of dust can be revealed.  Dwarf galaxies of the Local Group are {\\revised nearby} laboratories to spatially study the life cycle of the different dust components and the metal enrichment of low-metallicity ISM. NGC6822 is our closest \\citep[490 kpc;][]{Mateo1998} metal-poor galaxy neighbour ($\\sim$ 30$\\%$ \\zsun) beyond the Magellanic Clouds and possesses isolated star forming (SF) regions at different evolutionary stages. The galaxy is a perfect candidate to study the feedback of the star formation on the low-metallicity ISM by analysing the spatial variations of its SEDs. NGC6822 also possesses an intriguing rotating H~{\\sc i} disk of 1.34 $\\times$ 10$^{8}$ \\msun, that extends far beyond the optical disk \\citep{Mateo1998} and has one of the largest HI holes ever observed in a dwarf galaxy \\citep{deBlok2000}.  NGC6822 was observed in 2009 October as part of the Science Demonstration observations for the Dwarf Galaxy Survey (PI :  S. Madden), with the instruments PACS and SPIRE at 70, 100, 160 and 250, 350 and 500 \\mic\\ respectively. At SPIRE 500 \\mic\\ (36\"), we can resolve ISM structures of $\\sim$ 85pc, spatially sufficient to accurately probe the distribution of dust temperature throughout the galaxy, especially dust in its coldest phases. ",
        "conclusions": "We present \\hersc\\ images of NGC6822 which resolve ISM structures up to 500 \\mic. We find that the 250/500 ratio (tracing the cold dust temperature range) may be dependent on the 24 \\mic\\ surface brightness and thus {\\revised trace the} SF activity. We model individual SEDs across NGC6822 and show that the SED shape is evolving from H~{\\sc ii} regions to less active regions, with H~{\\sc ii} regions having a {\\revised warmer dust temperature} range. We derive very high dust masses using graphite to describe carbon dust and find that the use of amorphous carbon decreases the dust masses, indicating that SED models including \\hersc\\ constraints require different dust properties, namely more emissive grains."
    },
    "1005/1005.3735_arXiv.txt": {
        "abstract": "We consider the dilaton in the strong string coupling limit  and elaborate on the original idea of Damour and Polyakov whereby the dilaton coupling to matter has a minimum with a vanishing value at finite field-value. Combining this type of coupling with an exponential potential, the effective potential of the dilaton becomes matter density dependent. We study the background cosmology, showing that the dilaton can play the role of dark energy. We also analyse the constraints imposed by the absence of violation of the equivalence principle. Imposing these constraints and assuming that the dilaton plays the role of dark energy, we consider the consequences of the dilaton on large scale structures and in particular the behaviour of the slip functions and the growth index at low redshift. ",
        "introduction": "The observed late time acceleration of the Universe has no clear theoretical explanation yet. One of the putative candidates is dark energy whereby a scalar field with a flat enough potential provides the potential energy leading to the accelerating phase (for reviews and references see e.g. \\cite{various}). Although this seems like a natural scenario, it is fraught with difficulties. The most important one is certainly the absence of a complete understanding of the role of quantum corrections in such models. This is akin to the hierarchy problem of Higgs physics, albeit even more serious due to the stringent phenomenological constraints that acceleration imposes on the dark energy potential. In particular, both the vacuum energy and the mass of the scalar field must be minute. The former is nothing but a reformulation of the cosmological constant problem whereas the latter prescribes that there should exist a new fifth force complementing the gravitational interaction at very large scales. Reconciling such a long range force with local experiments of gravity on earth and in the solar system is a difficult task. Three types of mechanisms can be generically invoked. The first one appears in the DGP \\cite{DGP} modification of gravity where the Vainshtein \\cite{Vainshtein} effect is present locally. In these models, gravity is modified on large scales and preserved close to massive bodies due to the non-linearities of the scalar field kinetic terms. The shielding of the scalar field by massive bodies is also a feature of chameleon models \\cite{chameleon1,chameleon2,chameleon3} (and therefore of $f(R)$-gravity \\cite{f(R)chameleon}) where the mass of the scalar field becomes environmentally dependent. This leads to a thin shell effect preventing any deviations from Newton's law in the vicinity of massive objects.  Another mechanism has been advocated in a string theoretic context: the Damour-Polyakov effect \\cite{Damour:1994zq}. Considering the string dilaton in the strong coupling regime, it turns out that no violation of general relativity would be observed provided the coupling of the dilaton to matter were driven to zero by the cosmological expansion. This results stands when no potential is taken into account for the dilaton. It turns out that in the strong coupling regime, one expects that the dilaton will have an exponentially decreasing potential akin to the ones used to describe dark energy. This was noticed in \\cite{ven1, ven2, ven3} and \\cite{ven1} proposed a scenario whereby quintessence was the runaway dilaton, though the couplings to matter were neglected here. Later this was modified and  gravitational tests evaded provided the coupling of the dilaton to matter vanishes for an infinitely large dilaton. This is equivalent  to the Damour-Polyakov mechanism with a minimum at infinity.  In this paper, we will focus on the original Damour-Polyakov setting where the coupling vanishes for a finite value of the dilaton while keeping an exponentially runaway dilaton potential. In this case, the potential term tends to displace the dilaton from the minimum with no coupling to matter. The fifth force constraints would not be evaded anymore. In fact this result only stands when the string and Planck scales are of the same order of magnitude. Provided the string scale is lower that the Planck scale by a few orders of magnitude, we find that the Damour-Polyakov mechanism is at play albeit only locally where matter densities are large. This allows one to evade solar system constraints on gravity. This environmentally dependent Damour-Polyakov mechanism implies that there exists a fifth force whose manifestation is prominent on galaxy cluster scales. This prompts the interesting possibility that one could probe and measure dilatonic modifications of gravity with future large scale surveys.  This paper is organised as follows. In Section 2, we will recall the main properties of dilaton models in the strong coupling regime. Then, in Section 3 we will study the local tests and the constraints they impose on dilatonic models and investigate the chameleon mechanism as well as the Damour-Polyakov mechanism. In Section 4, we present the environmentally dependent Damour-Polyakov mechanism and discuss the cosmological evolution as well as the the local constraints. In Section 5 we will focus on the  consequences for large scale structures in the universe. Our conclusions can be found in Section 6. ",
        "conclusions": "We have presented new results on dilaton models in the strong coupling regime. Our analysis combines a runaway potential as suggested in \\cite{ven1} with a coupling to matter with a vanishing minimum, as investigated in \\cite{Damour:1994zq}.  In the absence of any potential term, the coupling to matter is enough to drive the dilaton to its minimum cosmologically \\cite{Damour:1994zq}. As a result, the dilaton would evade gravitational tests on the non-existence of fifth forces. This result is jeopardised by the runaway dilatonic potential, in particular if the dilaton plays the role of dark energy. In this paper we have focussed on the dilaton as a candidate for dark energy and shown how this can be compatible with gravitational tests of gravity, but could lead to observational tests at large scales.  We have shown that for models where the string scale is lower than the Planck scale, the Damour-Polyakov mechanism, whereby the coupling to matter vanishes dynamically, is at play only locally in the presence of large enough over-densities such as the one present in galaxies. On larger cosmological scales, the fifth force is not suppressed implying the presence of a significant modification of gravity. We found that this effect could be relevant on galaxy cluster scales where the growth of structures would be affected. As a result, the future galaxy survey could give stringent constraints on dilaton models and have the potential to indicate the possible existence of a dilaton in the strong coupling regime. We also expect that the non-linear growth of structures would also be affected by the presence of a dilaton. This is left for future work."
    },
    "1005/1005.1506_arXiv.txt": {
        "abstract": "{ Infrared dark clouds (IRDCs) are cold and dense reservoirs of gas potentially available to form stars. Many of these clouds are likely to be pristine structures representing the initial conditions for star formation. The study presented here aims to construct and analyze accurate column density and dust temperature maps of IRDCs by using the first Herschel data from the Hi-GAL galactic plane survey. These fundamental quantities, are essential for understanding  processes such as fragmentation in the early stages of the formation of stars in molecular clouds. We have developed a simple pixel-by-pixel SED fitting method, which  accounts for the  background emission. By fitting a grey-body function at each position, we recover the spatial variations  in both the dust column density and temperature  within the IRDCs. This method is applied to a sample of 22 IRDCs exhibiting a range of angular sizes and peak column densities. Our analysis shows that the dust temperature decreases significantly within IRDCs, from background temperatures of 20-30~K  to minimum temperatures of 8-15~K within the clouds, showing that dense molecular clouds are not isothermal. Temperature gradients have most likely an important impact on the fragmentation of IRDCs. Local temperature minima are strongly correlated  with column density peaks, which in a few cases reach N$_{H_2} = 1\\times 10^{23}$~cm$^{-2}$,  identifying these clouds as candidate massive prestellar cores. Applying this technique to the full Hi-GAL data set  will provide important constraints on the fragmentation and thermal properties of IRDCs, and help identify hundreds of massive prestellar core candidates. } ",
        "introduction": "One of the major issues in star formation is understanding the physical conditions, including the temperature and density structure, of the material out of which stars will form.  While fundamental, these quantities are difficult to determine. These regions are cold and therefore radiate energy away through dust at the far infrared/submillimeter wavelengths, a frequency range barely accessible from the ground because of the opacity of the atmosphere.  For the first time, a high angular resolution view of the Universe from 70 to 500 \\microns\\ is now becoming available thanks to the ESA Herschel Space Observatory \\citep{pilbratt2010}. The frequency range provided by its two photometry instruments PACS \\citep{poglitsch2010} and SPIRE \\citep{griffin2010} perfectly matches the emission peaks of young stellar objects and prestellar cores, the direct progenitors of stars. It is  now possible to construct well constrained SEDs for these objects, and therefore determine the spatial distributions of the temperature and column density towards a wide range of different objects.  Hi-GAL is the open time key project  to observe the inner part ($|l|\\leq60\\degr; |b| \\leq1\\degr$) of the Galactic plane with Herschel \\citep{molinari2010a}.  The sensitivity of this survey will allow the detection of 10~M$_{\\odot}$ molecular clouds at a distance of 4~kpc. A large number of diverse studies will  result from this survey covering scientific themes such as evolved stars, diffuse interstellar medium, and star formation. In this  Letter, we present the first  Herschel study analyzing the temperature and column density structure of infrared dark clouds (IRDCs). These objects are reservoirs of cold, dense molecular gas seen in extinction in the mid-infrared against the strong background emission of the Galactic plane \\citep[Fig.~\\ref{sdc2col}]{perault1996,rathborne2006}. This gas is only slightly processed by current star formation activity and most likely still contains the initial conditions of star formation imprinted within them. Their study can therefore illuminate the earliest stages of star formation. In Sect. 2 of this paper, we describe the data, Sect. 3 focuses on the background properties. Section 4 presents the first results, while Sects. 5 and 6 summarize the discussion and conclusions of the paper, respectively. ",
        "conclusions": "We have developed a technique to obtain reliable temperature and column density maps of IRDCs using the first Herschel data from the Hi-GAL open-time key project. We have demonstrated our method for a small sample of massive IRDCs by showing that they are non-isothermal but have a definite temperature structure, their inner regions being cooler than their edges. This temperature structure probably has an important impact on the fragmentation of the cloud in the early stages of its evolution. We also identified  a number of high column density peaks, reaching values of $10^{23}$~cm$^{-2}$, which are most likely to be the birth sites of massive stars. The full Hi-GAL dataset is expected to contain several hundreds of these sources.     \\begin{figure}[!t!] \\hspace{-0.cm} \\vspace{-0.0cm} \\includegraphics[width=8.5cm,angle=0]{temp_coldens.pdf} \\caption{Dust temperature versus column density plot. The blue square symbols show the average column density and mass weighted temperature for the 22 IRDCs analyzed. The red circular symbols show the temperature and column density at the column density peaks within the IRDCs. The filled symbols correspond to IRDCs with at least one 70~\\microns point-like source, i.e., star-forming IRDC/peaks. \\label{TvsN}} \\end{figure}"
    },
    "1005/1005.3833.txt": {
        "abstract": "\\label{sec:Introd}  During the 1970s members of the British Interplanetary Society embarked on a landmark theoretical engineering design study to send a probe to Barnard's star. Project Daedalus was a two-stage vehicle employing electron beam driven inertial confinement fusion engines to reach its target destination. This paper sets out the proposal for a successor interstellar design study called Project Icarus. This is an attempt to redesign the Daedalus vehicle with similar terms of reference. The aim of this study is to evolve an improved engineering design and move us closer to achieving interstellar exploration. Although this paper does not discuss prematurely what design modification are likely to occur some indications are given from the nature of the discussions. This paper is a submission of the Project Icarus Study Group. ",
        "introduction": "Sending robotic probes to solar systems other than our own is a vastly different technical challenge when compared against interplanetary exploration. The nearest stars are over four light years away or a distance of $\\approx272,000$ Astronomical Units (AU). For this reason, history has shown a reluctance to accept the prospects for interstellar travel as realistic. Since 1999 however, astronomical observations of other stars have identified over 400 extra-solar planets, primarily using the radial velocity method \\cite{Schneider,Butler,Udry}. This has led to the growing possibility that within decades astronomers may identify a habitable world which could, in theory, be one day explored by humans. The technical challenge remains however - even when you know where you want to go how do you get there?  There have been numerous theoretical studies into different propulsion schemes one may use to cross the vast distance of space. A quick review of the various concepts will immediately show that many of those schemes are ruled out. This includes chemical, electric and nuclear fission \\cite{Long1}. This is largely due to either a lack of specific impulse or a lack of thrust to give high exhaust velocity. A high exhaust velocity would be necessary to reach the speeds that are required to travel to another star within the time span of a human lifetime - one of the Terms of Reference of the Icarus mission to be discussed later in this paper. Fuels with comparatively low levels of energy available, for liberation in the form of thrust, would also require a high mass ratio placing constraints on how far the spacecraft can go. \\cite{Mallove}  Although there are exotic concepts for interstellar travel such as warp drive and wormholes, these are currently considered to be purely speculative \\cite{Long2,Obousy} . Concepts such as antimatter have been investigated but the technology is currently too immature to harness for a spacecraft \\cite{Forward1}. Employing the energy from the Sun in solar sail driven vehicles is credible but has the problem that solar intensity reduces inversely with the distance squared. This can be compensated for by using large collimated laser beams \\cite{Forward2}, but this technology has not been demonstrated for such an application. In the search for fuels which are energetic, provide for low mass ratios, low Thrust/weight ratio for high exhaust velocities, designers are led to consider nuclear pulse engines, in particular with a fusion based fuel.  The idea of using nuclear pulse propulsion was first proposed by Stanislaw Ulam in 1947 and subsequently resulted in an engineering study led by Ted Taylor \\cite{Taylor}. This was to employ nuclear pulse technology to propel a vehicle by capturing the blast products of a nuclear explosion on the rear of a giant pusher plate, transferring momentum to the spacecraft and its occupants, cushioned from the blast by several shock absorbers. This eventually led to the Project Orion engineering study with around \\$11 million being spent over a period of 7 years. The design was for a mission which would take only $approx$140 years to reach Alpha Centauri with 1 `unit of propellant' being exploded every 3 seconds to push the spacecraft at 1g acceleration to $\\approx3\\%$ of light speed [11]. However, with the arrival of various test ban treaties external nuclear pulse was shelved to the archives. Despite this, Project Orion remains the first comprehensive interstellar spacecraft design which was credible.  In 1971 Friedwardt Winterberg explored the concept of using Marx generators to power electron particle beams \\cite{Winterberg}. This idea was picked up by members of the British Interplanetary Society (BIS) who were considering embarking on an engineering design study to demonstrate that interstellar travel was, at least, possible with current, or near future, technology. Then in January 1973 members of the BIS first met to discuss the challenges of interstellar propulsion and the idea of Project Daedalus was born. Led by Alan Bond, Tony Martin and Bob Parkinson, members came together to create what has become one of the most comprehensive interstellar engineering studies ever undertaken. Project Daedalus had three stated guidelines:  \\begin{enumerate} \\item The spacecraft must use current or near future technology. \\item\tThe spacecraft must reach its destination within a human lifetime. \\item\tThe spacecraft must be designed to allow for a variety of target stars. \\end{enumerate}  The members of the Daedalus study group were all volunteers but with a solid knowledge of engineering and science. The final design was published in 1978 \\cite{Bond} and was a two-stage spacecraft nearly $200$m in length powered by electron driven D/He$^3$ fusion reactions, eventually accelerating up to $12$\\% of light speed to arrive at its target destination Barnard\u0092s star 5.9 light years away in under $50$ years. The Daedalus first stage had a structural mass of $1,690$ tonnes with $46,000$ tonnes of propellant in six tanks and it would burn for 2.05 years before jettisoning it. The second stage had a structural mass of 980 tonnes with 4,000 tonnes of propellant in four smaller tanks and would burn for 1.76 years, the tanks all being jettisoned prior to reaching the destination. Both stages would achieve an exhaust velocity of around 10,000 km/s from the detonation of the D/He$^3$ pellets at a frequency of 250 Hz. The second stage carried a $450$ tonne scientific payload with 18 probes to be used to study the target solar system planets. One of the interesting problems that Daedalus had to deal with is the transmission of a radio signal over vast distances back to earth. The team proposed the unique solution of using the 2nd stage parabolic reaction chamber to achieve high antenna gain. It is too early in the design process (pre-concept stage) to identify how Project Icarus may do things differently, giving the changes in technology over the three decades. However, at this stage we can speculate that likely design modifications will be in the areas of electronics, computing science probes, pellet ignition driver, the propellant and its acquisition. Some of the issues will be briefly discussed in this paper.  Project Daedalus also addressed the fundamental question of the whereabouts of intelligent life in the universe given that sufficient time had existed for sun-like stars to evolve with potentially habitable planets, despite the large distances between the stars, a problem termed the Fermi Paradox. Project Daedalus showed that if humans could envisage a vehicle capable of crossing those distances in a reasonable amount of time, then later more efficient ways of doing it would be found. The project showed that interstellar flight was most likely possible and more sophisticated thinking about the absence of intelligent visitors would be required \\cite{Parkinson}.  In addition to Project Daedalus, a number of alternative fusion based propulsion studies have also been performed. One of these was Project Vista \\cite{Orth} in the 1980's which proposed using D/T fusion reactions driven by a $\\approx 5$ MJ laser beam to execute a mission to Mars within 6 months or any planet in the solar system within 7 years. Vista would employ a technique known as fast ignition to achieve an energy gain compared to the input energy of $>1000$ and a specific impulse of $\\approx10^4 {\\rm s}$. Another study in the 1980's was Project Longshot \\cite{Beals} which proposed using D/He$^3$ reactions driven by either laser or electron beams to accelerate a mission to Alpha Centauri within a century by reaching a speed of $5\\%$ of light and by having a specific impulse of $10^6 {\\rm s}$. The He$^3$ was to be manufactured on earth using particle accelerators which bombard a lithium blanket. Alternatively, it was to be captured from the solar wind or mined en route from the atmosphere of Jupiter.  The most distant robotic pioneers sent out of the solar system are the Pioneer and Voyager spacecraft, launched in the 1970's. Travelling at slow speeds of only a few AU per year, these spacecraft would take tens of thousands of years to reach the distance of nearest stars if they continued in their current direction. There is a large technological gap which needs to be breached, if interstellar exploration is to become a reality. How do we bridge this gap? One answer is to perform a focussed design study for a particular propulsion engine with the end objective of increasing the Technology Readiness Level or TRL of the design type. These describe a range of technological maturities from conjecture through to a tested application \\cite{Mankins}.  Hence one of the main purposes behind Project Icarus and the decision to base the engine design on mainly fusion technology. ",
        "conclusions": "It has been over three decades since the landmark Daedalus engineering study and Project Icarus has finally arrived. This will be a complete redesign of the Daedalus systems including a re-examination of some of the original assumptions. An international team is now assembling to work on this exciting endeavour and bring the human dream of interstellar travel that much closer to being possible. Time will show if this ambitious project meets its overall goals in advancing fusion based space propulsion.  This work has been supported by the project consultants Marc Millis, Paul Gilster, Greg Matloff and Tibor Pacher in a personal capacity. Further information on Project Icarus is available from the online web site (www.icarusinterstellar.org). Since the submission of this paper, several new additions have joined the Project Icarus Study Group, some of which were present at the symposium. This includes Ian Crawford, Pat Galea, Adam Crowl, Robert Adams, Rob Swinney, Andreas Hein, Philip Reiss, Jardine-Barrington-Cook and James French. \\\\ \\\\ {\\it Of Icarus, In ancient days two aviators procured to themselves wings. Daedalus flew safely through the middle air and was duly honoured on his landing. Icarus soared upwards to the sun till the wax melted which bound his wings and his flight ended in fiasco. The classical authorities tell us, of course, that he was only ``doing a stunt''; but I prefer to think of him as the man who brought to light a serious constructional defect in the flying-machines of his day. So, too, in science. Cautious Daedalus will apply his theories where he feels confident they will safely go; but by his excess of caution their hidden weaknesses remain undiscovered. Icarus will strain his theories to the breaking-point till the weak joints gape. For the mere adventure? Perhaps partly, this is human nature. But if he is destined not yet to reach the sun and solve finally the riddle of its construction, we may at least hope to learn from his journey some hints to build a better machine} \\cite{Eddington}."
    },
    "1005/1005.3029_arXiv.txt": {
        "abstract": "On 2010 March 19, the {\\it Swift}/Burst Alert Telescope triggered on a short burst with temporal and spectral characteristics similar to those of Soft Gamma Repeater (SGR) bursts. The source location, however, did not coincide with any known SGR. Subsequent observations of the source error box with the {\\it Swift}/X-ray Telescope and the Rossi X-ray Timing Explorer ({\\it RXTE}) led to the discovery of a new X-ray source, with a spin period of 7.56 s, confirming \\sgr as a new magnetar. Based on our detailed temporal and spectral analyses, we show that the new SGR is rapidly spinning down (4$\\times$10$^{-12}$ s/s) and find an inferred dipole magnetic field of 1.8$\\times$10$^{14}$ G. We also show that the X-ray flux of \\sgr remained constant for approximately 20 days following the burst and then started to decline. We derived an accurate location of the source with the {\\it Chandra} X-ray Observatory and we searched for a counterpart in deep optical and infrared observations of \\sgrnos, and for radio pulsed emission with the Westerbork Radio Synthesis Telescope. Finally, we compare the spectral and temporal properties of the source to the other magnetar candidates. ",
        "introduction": "The first confirmed Soft Gamma Repeater (SGR) burst was detected in 1979 but it took $\\sim7$ years (\\citealt{laros87,atteia87,kouv87}) to associate it with a newly discovered population of astronomical objects, namely neutron stars repeatedly emitting soft and short gamma-ray bursts with peculiar properties. A decade later, these objects, together with their cousins, Anomalous X-ray Pulsars (AXPs; see e.g., \\citealt{mer08}), were shown to possess extremely high magnetic fields ($\\sim10^{15}$ G), as deduced from their spin periods and spindown rates, and were named magnetars \\citep{dt92, td96, kouv98, vas97}.  The most distinguishing magnetar characteristics are a narrow range of spin periods (2$-$12s), high spindown rates ($\\sim10^{-10}$ s/s) and unpredictable burst-active episodes. The latter can have a multitude of bursts with a wide luminosity range, anywhere from 10$^{-2}$ to 10$^7$ times the Eddington luminosity. It is these bursts that enable the initial classification of a source as an SGR, to be subsequently confirmed as a magnetar with timing analysis (see \\citealt{woo06} and \\citealt{mer08} for a detailed review on magnetars.)  By 2008 a dozen sources were confirmed magnetars and a couple more were listed as magnetar candidates pending observations of new burst active episodes. For the known sources, frequent spin period monitoring observations with the {\\it Rossi} X-ray Timing Explorer ({\\it RXTE}) over the last fifteen years have provided a wealth of insight about their dynamics and energetics. Interestingly, despite their common characteristics, there are several important differences across this small population, and a larger sample of objects at different evolutionary stages is necessary to better understand the phenomenon. Magnetar hunting however, is not an easy task; we have two strikes against us {\\it ab initio}:  their very small numbers and their very short active periods (days to months). Thus, the best way to discover magnetars is the combination of an all-sky X-and gamma-ray monitor -- to detect their bursts -- and a timing instrument -- to measure their spin and spindown rates, and estimate their magnetic fields.  On 2010 March 19, the Burst Alert Telescope (BAT) onboard NASA's  {\\it Swift} satellite triggered at 18:34:50 UT on one short and soft burst (Barthelmy et al. 2010) originating from a region in the plane of our Milky Way galaxy that harbors at least three other magnetar candidates: AXP XTE J1810-197, and SGRs 1806-20 and 1801-23. Using the initial source localization with the BAT, which had a  3$\\arcmin$ error radius, we were able to excluded all these sources; further, based on the spectral and temporal properties of the event, we identified the source of the emission as a possible new SGR. Subsequent observations (starting from 60~s after the BAT trigger) of the source region with the {\\it Swift}/X-Ray Telescope (XRT) provided the source position with a 1.6$\\arcsec$ uncertainty \\citep{gel10}. {\\it Swift}/Ultraviolet Optical Telescope (UVOT) observations starting 71 s after the trigger did not detect a source in the UV or optical bandpass at the XRT position \\citep{margel10}. To confirm the nature of the source, we observed it with {\\it RXTE} for 1.3 ks {\\it less than 3.25 hours} after its discovery was reported. A subsequent timing analysis of this observation revealed coherent pulsations in the X-ray flux of the source with a period of 7.56 s \\citep{gogus10a}. Thus, we were able to identify \\sgr as a slowly rotating neutron star in less than 6 hours after its discovery, the second time such a fast discovery was enabled only due to the unprecedented synergy of {\\it Swift} and {\\it RXTE} (the first was the identification of SGR J0501+4516; \\citealt{gogus09}).  {\\it Swift} was the only satellite that triggered on \\sgr; although the source was in the field of view of the Gamma ray Burst Monitor (GBM) onboard the Fermi Gamma-ray Space Telescope ({\\it FGST}), the fluence of the event was much lower than any of the GBM trigger thresholds. To precisely locate the source we triggered our Target of Opportunity (ToO) observations with the {\\it Chandra} X-Ray Observatory ({\\it CXO}) and obtained observations with the European Southern Observatory ({\\it ESO}) Very Large Telescope/High Accuity Wide-field K-band Imager ({\\it VLT/HAWKI}) and with the Gran Telescopio Canarias/Optical Imager and Spectrograph ({\\it GTC/OSIRIS}). We also observed with the Westerbork Synthesis Radio Telescope Pulsar Machine II ({\\it WSRT/PuMa II}) and searched for radio pulsations from the source.  In this paper, we report on the discovery of \\sgrnos, and introduce its characteristics through our detailed multi-wavelenth investigations. In \\S 2, we describe all multi-wavelength data used in our study. We present in \\S 3 the precise localization of the new source using {\\it CXO}, 2MASS and UKIRT data, and the temporal and spectral characteristics of the persistent X-ray emission from the source using our {\\it Swift} and {\\it RXTE} observations. Finally, we discuss and interpret our results in \\S 4. ",
        "conclusions": "\\subsection{Burst emission}  At the onset of its active episode, \\sgr emitted a single, relatively faint, burst that was seen with {\\it Swift}/BAT and INTEGRAL \\citep{kui10}; the total event fluence (15$-$350 keV) was 1.3$\\times$10$^{-8}$ erg/cm$^2$. Compared to the persistent emission flux level (2$-$10 keV), the burst energetics are very low: it would take only 13 minutes for the persistent source to emit an equivalent amount of energy. It is, therefore, unlikely that the (detected) burst emission provided the energy budget for the persistent X-ray emission.  The spectrum of the {\\it Swift}/BAT burst is well described by a power law with exponential cutoff (E$_{peak}$ = 38 keV) or a black body with kT=10 keV. Its spectral and temporal characteristics are very similar to those of known magnetars, in particular, to those of SGR~J0418$+$5729, which emitted only two rather dim bursts in one day \\citep{van10}, and 1E 2259$+$586, which displayed a short episode of bursting for about 10 ks in 2002 seen only in hard X-rays \\citep{kas03}. In addition to the {\\it Swift}/BAT event, we identified 4 weak X-ray bursts in our monitoring observations with the {\\it RXTE}/PCA, which were not detected in the simultaneous {\\it Fermi}/GBM data. It is interesting to note that the {\\it RXTE}/PCA bursts are very similar in duration and X-ray intensity to the five short and rather weak bursts, also detected in the {\\it RXTE}/PCA passband, emitted by the young, relatively strongly magnetized (B$_{dipole}$=4.9$\\times$10$^{13}$ G) rotation powered pulsar PSR J1846$-$0258 \\citep{gav09}.   \\subsection{Persistent emission properties}  The spin period, 7.5654091 s,  and period derivative, 4.39$\\times10^{-12}$ s/s, of the source derived using the {\\it RXTE} data are consistent (within 1.5$\\sigma$) with the values obtained by \\citet{esp10}. Using these values, we estimate the dipole magnetic field strength of \\sgr to be  B = 1.8 $\\times$10$^{14}$ G, which is at the lower end of the B-fields of SGRs and close to the median value of those of AXPs (B$_{\\rm median}$ = 2.2$\\times$10$^{14}$ G). Note that the two recently discovered SGRs J$0418+5729$ and J$0501+4516$ possess the lowest SGR surface magnetic fields of $<$ 3$\\times$10$^{13}$ G and 1.7$\\times$10$^{14}$ G, respectively.  Figure \\ref{fig:xrtflux} exhibits the evolution of the persistent X-ray emission unabsorbed fluxes ($1-10$ keV) for five magnetars (data obtained from \\citealt{rea09}) and \\sgrnos, observed with {\\it Swift}/XRT in the $2-10$ keV range. Each source was observed immediately after an outburst, when the persistent emission level had significantly increased from its quiescent state. We notice that all sources had similar maxima at the onset of their activation, ranging between $10^{-11}$ and $10^{-10}$ erg/cm$^2$/s. Three of the sources (SGRs~J$1627-41$, J$0501+4516$ and J$1900+14$) declined in a similar monotonic manner. However, SGR~J$1627-41$ settled at the lowest flux level of all sources ($\\sim10^{-12}$ erg/cm$^2$/s), while SGR~J$1900+14$ reached a flat plateau an order of magnitude higher and SGR~J$0501+4516$ was still declining at the cessation of the XRT observations. The other two sources (1E~$1547-5408$ and CXOU~J$164710-455216$) show similar trends (but different from the other three sources): they remain flat after activation for $10-20$ days and decline {\\it thereafter} with the same trend as the former three.  \\begin{figure}[h] \\vspace{0.0in} \\centerline{ \\includegraphics[angle=-90,width=6.0in]{f6.eps}} \\caption{Unabsorbed flux of \\sgr (filled circles) obtained by fitting power law (upper values) and blackbody (lower values) models to XRT data in PC mode, compared with other sources [SGR 1900+14 (plus), 1E 1547.0$-$5408 (triangles), SGR 0501+4516 (cross), CXOU J164710$-$455216 (diamonds) and SGR 1627$-$41 (squares); data obtained from \\citealt{rea09}]. Time is measured from the onset of burst active episodes. The horizontal dashed line indicates the {\\it Chandra} 2$\\sigma$ flux upper limit of \\sgr from 2009 February 13.} \\label{fig:xrtflux} \\end{figure}  We have plotted the \\sgr data as a band with an upper and lower bound on Figure \\ref{fig:xrtflux}; the upper flux limit is derived when we fit a single power law and the lower with a blackbody fit. Because of the excessive column density towards the source, there are very few photons below 2 keV and the spectral fits are not well constrained; the real unabsorbed flux of the source is within these two bounds, and seems to remain constant over the first $\\sim$20 days following the onset. Interestingly, we found an archival {\\it Chandra}/ACIS-S observation covering the field of \\sgr starting on 2009 February 13 10:50:20 UT for a total of 8 ks. The source was not detected within our \\cxo error circle (\\S\\ref{sec:chandraobs}) and we were able to derive a $2\\sigma$ upper limit (indicated by a dotted line on Figure \\ref{fig:xrtflux}) for its quiescent unabsorbed flux of $3.4\\times 10^{-13}$ erg/cm$^2$/s \\citep{gogus10b}. This is almost 100 times fainter than the current state of the source; assuming that \\sgr reaches this level again, it would take at least 2 years if it follows the same decline trend of the other five sources as indicated by the last XRT observations. The source quiescent level is close the lowest flux level of SGR 1627$-$41, the faintest magnetar thus far (Mereghetti et al. 2006, Esposito et al. 2009), and also to the flux of the faintest dim isolated neutron star RX J0420.0$-$5022 (Haberl et al. 2004) of 4.8 $\\times$ 10$^{-13}$ erg/cm$^2$/s.  Finally, pulsed radio emission has been detected from two magnetars, XTE\\,J1810-197 \\citep{camilo2006} and 1E\\,1547.0-5408 \\citep[SGR\\,J1550-5418;][]{camilo2007}. The former showed radio pulsations after an X-ray outburst from the source, but the latter was seen in radio while the source was quiescent in X-rays. Searches for pulsed radio emission from other magnetars have been unsuccessful \\citep[e.g.][]{burgay2006,crawford2007,hessels2008}. Our search for radio emission from \\sgr with the {\\it WSRT} resulted in non-detections, with upper limits of $0.2-0.6$ and $0.1-0.3$,mJy (depending on the assumed pulse widths) at 2280 and 1380\\,MHz, respectively. \\citet{burgay2010} reported an upper limit of 0.09 mJy at 1400 MHz on 2010 March 25 with the Parkes radio telescope. These upper limits are somewhat higher than some of the deepest radio searches for other magnetars, but more than an order of magnitude lower than the detections of XTE\\,J$1810-197$ and 1E\\,$1547.0-5408$.  \\subsection{Environment of \\sgr}  A visual inspection of the region surrounding \\sgr on the $K_{\\rm s}-$band images of the {\\it UKIRT~} GPS reveals no obvious source clustering in the vicinity of the source. However, as the field is small and crowded, low mass open clusters may be difficult to discern from bright background sources. We have performed a color-color search to uncover any cluster members; we caution, however, that it is still likely that the number of ``unrelated'' sources may obscure the detection of any faint, low surface brightness cluster that may be present.  There are several sources in the larger field which are plausible sites for the birth of \\sgrnos. One of these is the HII region IRAS~$18310-0825$, at $RA=18^{\\rm h}33^{\\rm m}47^{\\rm s}.9$, $Dec = -08^\\circ23\\arcmin52\\arcsec$, which is offset by roughly $7\\arcmin$ from \\sgrnos, another is [LPH96] 023.162+0.023 ($RA = 18^{\\rm h}33^{\\rm m}23^{\\rm s}.9, Dec=-08^\\circ40\\arcmin 32\\arcsec$; 10 arcmin offset, \\citealt{loc96}), while a third is [KC97c] G023.5+00.0 at $RA= 18^{\\rm h}34^{\\rm m}20^{\\rm s}.0, Dec=-08^\\circ 21\\arcmin23\\arcsec.7$, which is also associated with a cluster known as BDS2003 118 \\citep{loc86}. Further, there are several dark clouds in the vicinity of \\sgrnos, which apparently do not yet host stellar populations. They have masses of $10^2-10^3$ M$_{\\odot}$, and are likely locations for a next generation of star formation within the region. Finally, there is also a known pulsar (PSR J1830-08) roughly 3'7 away from SGR J1833-0832 with an adopted distance of 5.67 kpc and a characteristic age of 147 kyr. It is possible that this source originates from the same star forming region associated with \\sgrnos.  While the dynamics of SGRs are as yet effectively unconstrained by observations (no SGRs have yet proper motion or radial velocity measurements) we can consider constraints by assuming that they i) have a short lifetime $<10^4$ years, and ii) have a kick velocity distribution broadly comparable to that seen for young pulsars. In practice this second assumption only improves the constraints we can make if the actual kicks are {\\em significantly} smaller, and hence limit the transverse distance that \\sgr could have travelled since birth. A typical pulsar kick might be of order 200 km s$^{-1}$, so that in $10^4$ years the source may move $\\sim 2$ pc on the sky. Hence the angular scale \\sgr could cover would be $\\theta \\approx 400\\arcsec / d_{kpc}$. For a distance of 6-12 kpc this corresponds to an angular scale of $30-70\\arcsec$, although allowing for a range of kick velocities might extend this range by a factor of 3 or more. In this case it would seem that none of the candidate clusters above are likely birthplaces. However, it is also possible (although not necessarily likely) that any progenitor may have ``evaporated'' from the cluster in which it formed with a velocity dispersion of a few km s$^{-1}$. With a lifetime of a few Myr, this may result in a progenitor at a large distance from the cluster, making our constraints weaker."
    },
    "1005/1005.5055_arXiv.txt": {
        "abstract": "{ As so far, the redshift of Gamma-ray bursts (GRBs) can extend to $z\\sim 8$ which makes it as a complementary probe of dark energy to supernova Ia (SN Ia). However, the calibration of GRBs is still a big challenge when they are used to constrain cosmological models. Though, the absolute magnitude of GRBs is still unknown, the slopes of GRBs correlations can be used as a useful constraint to dark energy in a completely cosmological model independent way. In this paper, we follow Wang's model-independent distance measurement method and calculate their values by using $109$ GRBs events via the so-called Amati relation. Then, we use the obtained model-independent distances to constrain $\\Lambda$CDM model as an example.}   \\arxivnumber{1005.5055}  \\begin{document} ",
        "introduction": "Since the discovery of an accelerated expansion of our Universe through the observations of supernova Ia (SN Ia) \\cite{ref:Riess98,ref:Perlmuter99}, many cosmic observations have been used to explore the nature of dark energy which has negative pressure and pushes the Universe into an accelerated phase. Particularly, the redshift of Gamma-ray burst (GRBs) can extend to higher redshift $z\\sim 8$. This merit makes it as a complementary cosmic probe to SN Ia. However, a big challenge, the so-called circular problem, has to be overcome when one calibrates the GRBs to obtain the distances at different redshifts. In \\cite{ref:Schaefer}, Schaefer {\\it et. al.} have derived the distance modulus in $\\Lambda$CDM model by using five GRBs correlations. The so-called circular problem would be committed  when one uses the resulted distance redshift relation to constrain any other cosmological models beyond $\\Lambda$CDM. To overcome this circular problem, Li, {\\it et. al} \\cite{ref:Lihong} put the GRBs correlation and cosmological model constraint together. Via the Markov Chain Monte Carlo (MCMC) method, they fixed the calibration parameters and constrained the model parameter space simultaneously. However, it looks like using a 'ruler' having no marks to measure the length of an object. And after the measurement, the length and scale of ruler are given together. So the lack of calibration makes the GRBs lose the power to constrain cosmological models. Cosmography method  was considered in \\cite{ref:cosmography} by parameterizing the luminosity distance $d_L$ in terms of deceleration $q_0$, jerk $j_0$ and snap $s_0$ parameters. In this way, the cosmological model dependent problem was removed. Liang {\\it et. al.} \\cite{ref:Liang} calibrated GRBs by using low redshift SN Ia and obtained a tight constrain to the cosmological model parameter space. This method was reconsidered by Wei \\cite{ref:Wei,ref:Wei109}. By analyzing the calibration process carefully, one can find the potential drawback. The calibrated GRBs correlation makes GRBs have the same luminosity distance-redshifts relation as SN Ia at lower redshifts. In this way, the relation is extended to higher redshifts. Equivalently, one just extends the luminosity distance-redshifts relation of SN Ia to higher redshifts. So, it makes the obtained luminosity distance-redshifts relation strongly depends on that of SN Ia. The worst thing is that it makes the data points of GRBs useless, because one has known the luminosity distance-redshifts relation at high redshifts from SN Ia, though no SN Ia is found at the high redshift regions of GRBs. Alternatively, Wang presented a model-independent distance measurement from GRBs calibrated internally \\cite{ref:wang}. The main point of Wang's method is that the statistical errors of correlation parameters $\\sigma_a$, $\\sigma_b$ and systematic error $\\sigma_{sys}$ obtained in $\\Lambda$CDM models are used, but are not the correlation parameters $a$ and $b$ themselves, for the definitions of $a$, $b$ please see Eq. (\\ref{eq:calib}). The viability of this implement comes from the observations that the errors of correlation parameters are almost the same for different values of $\\Omega_{m0}$ for $\\Lambda$CDM model, though the values of $a$ and $b$ are really different. Then, in terms of a set of model-independent distance measurements, the cosmic constraint from GRBs is set up via cubic spline interpolation from cosmological model independent distance ratio $\\bar{r}_p(z_i)$. The merits of this method are follows: (i) the constraint from GRBs is in a cosmological model independent way. It alleviates the circular problem. So, it can be used to constrain any other cosmological models. (ii) It is not calibrated by any other external data sets. It does not suffer any consistent problem when it is combined with other data sets as cosmic constraints. (iii) The cosmological model independent calibration is done firstly. It means that the 'ruler' has been marked. (iv) Though the absolute magnitude of GRBs is unkown, the slopes of GRBs correlations can be used as cosmological constraints.  Recently, Wei \\cite{ref:Wei109} used $109$ GRBs data points via Amati relation \\cite{r16,r17,r18} calibrated by SN Ia to constrain cosmological model. Based on the points mentioned above, to alleviate the data sets dependence and circular problem, we shall present our calculation results based on Wang's method via Amati correlation. In fact, we find a new $\\chi^2_{GRB}$ which only depends on the slopes of GRBs correlation and make the absolute magnitude of GRBs irrelevant. For the details, please see section \\ref{sec:method}. As results, a set of model-independent distance measurements are obtained on the basis of Amati relation that can be used to constrain cosmological models.  However we have to stress that the Amati relation has been criticized for many reasons in the literatures. Li demonstrated an ambiguity in determining the redshifts of GRBs \\cite{ref:Li}. Of course, this ambiguity can be overcome when the redshift is well determined. The major criticism came from Nakar and Prian \\cite{ref:Nakar} who developed a test for the Amati relation even in the case where the redshifts of GRBs were unknown. The test was also generalized by Band and Preece \\cite{ref:Band}. They concluded that the Amati relation suffered the problem of selection bias. Recently, the authors of \\cite{ref:Collazzi} have also concluded that the Amati relation is an artifact of selection effects with the burst population and the detector. They also point out that the Amati relation is failed whether or not the bursts have measured spectroscopic redshifts. If this is true for Amati relation, the results obtained based on Amati relation would be unreliable. But it is still in debating. So, before the dust settles down, we still assume that Amati relation is reliable in this paper.  This paper is structured as follows. In section \\ref{sec:method}, the values of cosmological parameter dependance in $\\Lambda$CDM model are presented as shown in \\cite{ref:wang}. The errors of correlation parameters $\\sigma_a$, $\\sigma_b$ and systematic error $\\sigma_{sys}$ are also calculated via Amati relation with $109$ GRBs data points. The $5$ bins model-independent distance measurements will also be found in this section. In section III, we use the resulted data points  to constrain $\\Lambda$CDM model as an example. A summary and discussion are put in section IV. ",
        "conclusions": "In this paper, by using $109$ GRBs data points via Amati relation, we have derived five data points of distance measurements which {\\it do not} depend on any cosmological models, i.e. in a model-independent way, based on the method firstly advocated by Wang \\cite{ref:wang}. Then it can be used to constrain any other cosmological models without the so-called circular and data sets consistence problem. We also find out that the concordance model ($\\Omega_{m0}=0.27$) is almost at the boundary of $1\\sigma$ regions of data points. So, one can reject the null hypothesis that Amati relation is based on cosmology slightly greater than $4.5\\sigma$ as shown in Fig. \\ref{fig:errorbars}.  When GRBs (the five data points of distance measurements) is used as a complementary cosmic constraint to $\\Lambda$CDM model as an example, the errors of model parameters are shrunken. Though the constraint is not much tighter than that calibrated via low redshift SN, the difference is very small. And, the results can be used to constrain other cosmological model beyond $\\Lambda$CDM. The important thing is that it {\\it does not} depend on SN Ia data set. The merits of this method has been mentioned in the introduction. Here, we list possible potential drawbacks of this completely model in-dependent method as follows: (i). The constraint is not much tighter than that obtained from calibration by using SN Ia data points. The looser may come from the lack of enough GRBs data points or the number of bins is not larger. (ii). How to design a cosmological model indicator, here it is model independent distance measurement $\\bar{r}_p(z)$, which is sensitive to distinguish the cosmological models may be the main points of this method. In this way, the lack of having no enough GRBs data points can be avoided in some senses. (iii) Another negative effect comes from the choice of fiducial  redshift $z_0$. Some average values of $r_p(z_0)$ would be better.  At last we have to alert the readers that our analysis is based on the assumption that Amati relation is reliable. However, in the literatures, the Amati relation has been criticized for many reasons as mentioned in the introduction. If that is true, our analysis would be unreliable."
    },
    "1005/1005.2675_arXiv.txt": {
        "abstract": "Quanz and coworkers have reported the discovery of the coolest known member of the Taurus star-forming complex (L2$\\pm0.5$) and Barrado and coworkers have identified a possible protostellar binary brown dwarf in the same region. We have performed infrared spectroscopy on the former and the brighter component of the latter to verify their substellar nature. The resulting spectra do not exhibit the strong steam absorption bands that are expected for cool objects, demonstrating that they are not young brown dwarfs. The optical magnitudes and colors for these sources are also indicative of background stars rather than members of Taurus. Although the fainter component of the candidate protostellar binary lacks spectroscopy, we conclude that it is a galaxy rather than a substellar member of Taurus based on its colors and the constraints on its proper motion. ",
        "introduction": "\\label{sec:intro}    Current surveys for brown dwarfs are exploring new regimes of temperature, mass, metallicity, age, binarity, and formation environment in order to address a variety of questions regarding objects at low masses and cold temperatures. For instance, brown dwarfs at low masses constrain the minimum mass of the initial mass function while the youngest brown dwarfs in the protostellar phase offer the most direct insight into the formation process of substellar objects.  The Taurus star-forming region ($\\tau\\sim1$~Myr, $d=140$~pc) has been the target of two recent surveys for low-mass brown dwarfs and protostellar brown dwarfs. \\citet{qua10} searched for the former by obtaining near-infrared (IR) images of a large sample of molecular cores in Taurus and combining these data with optical photometry from the Sloan Digital Sky Survey \\cite[SDSS;][]{yor00,fin04}. In the overlapping area of 1~deg$^{2}$ between the IR and optical images, \\citet{qua10} identified five new brown dwarf candidates and performed spectroscopy on two of them to assess their nature. Using these spectra, they classified one of the candidates, CAHA~Tau~1, as a young brown dwarf with a spectral type of L2, which would make it the coolest (and likely least massive) known member of Taurus \\citep{luh09tau}.  In another survey of Taurus, \\citet{bar09} used archival mid-IR images from the {\\it Spitzer Space Telescope} \\citep{wer04} to search for protostellar brown dwarfs. They reported the discovery of a $2\\farcs5$ pair of candidates that could comprise a binary system, which were designated as SSTB213 J041757.75+274105.5 A and B (henceforth J041757~A and B). Although spectroscopy was unavailable for these objects, \\citet{bar09} concluded that the proper motion of the A component supports its membership in Taurus and that the colors of the B component are inconsistent with a galaxy, which is the primary source of contamination in a survey for protostellar brown dwarfs.      We have performed IR spectroscopy on CAHA~Tau~1 and J041757~A to determine whether they are young brown dwarfs. We also have used the available colors and proper motion measurements for these two objects and the remaining candidates from \\citet{qua10} and \\citet{bar09} to determine if they are likely to be substellar members of Taurus. ",
        "conclusions": "We have used IR spectra, optical and IR colors, and proper motion measurements to examine whether CAHA~Tau~1--5 \\citep{qua10} and J041757~A and B \\citep{bar09} are substellar members of the Taurus star-forming region. We have obtained spectra of two of these sources, CAHA~Tau~1 and J041757~A. The absence of strong steam absorption in their spectra indicates that they are not young brown dwarfs. We classify CAHA~Tau~1 and J041757~A as reddened background stars with spectral types of M5--M6 and $\\lesssim$M2, respectively. All of the candidates from \\citet{qua10} and \\citet{bar09} appear below the sequence of known members of Taurus on an optical color-magnitude diagram, indicating that they are field stars or galaxies. One of these sources, CAHA~Tau~2, has $J-H$ and $I-K_s$ colors that are consistent with a late spectral type ($>$M6), but it is probably a field dwarf given its position in the optical color-magnitude diagram. Six of the candidates do not show excess emission from disks in their mid-IR colors. The remaining source, J041757~B, has a red spectral energy distribution that is consistent with both a protostar and an AGN. \\citet{bar09} concluded that this object is too red in $I-J$ and $J-[3.6]$ to be a galaxy, but the reddening of these colors relative to galaxies at high latitudes is easily explained by the Taurus dark cloud along its line of sight. The available proper motion measurements for CAHA~Tau~1--5 do not provide strong constraints on their membership in Taurus. We estimate a membership probability of 7\\% for J041757~A by comparing its proper motion to those of known Taurus members. The motion of this object is consistent with a field dwarf, which is how we classify it spectroscopically. The proper motion constraints for J041757~B are inconsistent with membership in Taurus and agree with the absence of motion expected for a galaxy. Based on these various data, we conclude that CAHA~Tau~1--5 and J041757~A and B are not substellar members of Taurus."
    },
    "1005/1005.1056_arXiv.txt": {
        "abstract": "We calculate IR divergent graviton one-loop corrections to scalar correlators in de Sitter space, and show that the leading IR contribution may be reproduced via simple semiclassical consistency relations.  One can likewise use such semiclassical relations to calculate leading IR corrections to correlators in slow-roll inflation.  The regulated corrections shift the tensor/scalar ratio and consistency relation of single field inflation, and non-gaussianity parameters averaged over very large distances.  For inflation of sufficient duration, for example arising from a chaotic inflationary scenario, these corrections become of order unity.  First-order corrections of this size indicate a breakdown of the perturbative expansion, and suggest the need for a non-perturbative description of the corresponding regime.  This is analogous to a situation argued to arise in black hole evolution, and to interfere with a sharp perturbative calculation of ``missing information\" in Hawking radiation. ",
        "introduction": "While inflation has become a very successful paradigm for early-universe physics, it still leaves us with many challenges and puzzles.  Perhaps the most important of these is incorporating it into a broader framework of quantum gravity, which allows systematic treatment of initial conditions, quantum corrections, {\\it etc.}  Even lacking such a treatment, quite a bit of detailed study can be accomplished, treating general relativity as an effective field theory and analyzing it perturbatively.  One of course faces the usual problem of nonrenormalizable UV behavior, but handles this by imagining that appropriate counterterms are provided by a more complete theory, and do not appreciably affect the long distance behavior of the theory relevant for inflationary predictions.  However, another puzzle arises in the infrared behavior of the theory:  various quantities in de Sitter space, or other inflationary cosmologies, exhibit  IR divergences.  IR divergences typically indicate that one has not understood an important aspect of the dynamics, or has asked a bad question.  For this reason these have been extensively discussed, with arguments ranging from the statement that they can be innocuously eliminated to the statement that they indicate a fundamental instability of de Sitter space causing the cosmological constant to relax to zero\\cite{Polyakov:1982ug,Polyakov:2007mm,Polyakov:2009nq,Tsamis:1992sx,Tsamis:1994ca,Tsamis:2007is,Myhrvold:1983hu}.  It is important to more clearly understand where such IR divergences occur, and in what way they modify calculation of physical quantities.  While IR divergent effects of massless scalars have long been known (see {\\it e.g.} \\cite{Linde:2005ht}), one might question how general these are, given that the correct theory may not have massless scalars.  However, in an effective description of gravity masslessness is inherent in a description in terms of gravitons.  This strongly motivates understanding the IR behavior of processes involving gravitational loops.  Loop calculations in gravity are notoriously complicated, due to the higher spin and its corresponding gauge invariance.  For this reason, so far results have been very limited.  One of the goals of this work is to improve this situation.  In particular, we calculate gravitational loop corrections to propagators in de Sitter space.  While we don't exhibit all such corrections, we note that some of the corrections are indeed IR divergent, and we extract this leading IR behavior.  Moreover, we argue that there is a simple way to derive such behavior from semiclassical arguments, related to the consistency relations of \\cite{Maldacena:2002vr,Creminelli:2004yq,Seery:2008ax} (and with antecedents in \\cite{Starobinsky:1994bd,Tsamis:2005hd}).  We explicitly check examples where these semiclassical relations match the detailed in-in loop calculations.  We then propose extending the semiclassical methods to more general, slow-roll, cosmologies, and show how they give leading IR corrections there.  We also demonstrate another check of these methods based on an explicit calculation.  Interestingly, we find that the corrections implied by these methods shift various cosmological observables, notably the tensor/scalar ratio and nongaussianity parameters in the ``large box\" (total inflated volume), and we exhibit  leading contributions to such shifts.  Since the corrections grow in the IR, for a sufficiently long period of inflation, for example in models with slow-roll after exit from a chaotic phase, the  loop corrections become of order one, compared to the tree-level quantities.  This suggests a breakdown of the perturbation series, associated with the strong IR effects.  This apparently sharpens the question of the physical interpretation of IR growth in our descriptions of inflationary spacetimes.  While there have been various suggestions that IR effects in de Sitter can be resummed, it seems possible that the breakdown of the perturbation series represents the need for new non-trivial non-perturbative dynamics.  Indeed, another motivation for this work is an interesting parallel, in black hole physics.  There are good indications that the extreme UV (superplanckian) regime of gravity is governed by physics classically approximated as black hole formation, providing a link between UV and IR physics.  A critical question, in light of Hawking's perturbative arguments\\cite{Hawking:1974sw} for information loss, is how such physics can be unitary.  Specifically, by a calculation of the state on a slice spanning the exterior and interior of a semiclassical black hole spacetime, one constructs an argument for information loss.  However, arguments were given in \\cite{Giddings:2007ie,Giddings:2009ae} that the perturbative state on such a slice ceases to be a good description, by a sufficiently long time, $t\\sim RS$, where $R$ and $S$ are the radius and the entropy of the black hole; one such argument was based on apparently large fluctuations in the perturbative state.  Failure of  such a sharp argument for information loss was proposed as a resolution of the information {\\it paradox}, leaving an information {\\it problem} of determining the non-perturbative gravitational mechanics that completes the description in this regime. Quantization on such slices in black hole spacetimes is very similar to quantization on spatial slices in inflationary spacetimes.  The calculations of this paper explicitly show perturbative corrections to the latter that become large on the time scale $t\\sim RS$, where now $R$ and $S$ are the de Sitter radius and entropy, in close parallel to the black hole story.\\footnote{This analogy was in particular pursued in ref.~\\cite{Nimatalk,ArkaniHamed:2007ky}, where it was pointed out than an attempt to regulate inflation by having it terminate through slow-roll fails if the termination time exceeds $t\\sim RS$.  The mechanism for this failure is the transition to self-reproduction.} This deepens the analogy between the situations, and amplifies the arguments for a similar breakdown in the black hole context.  In outline, the next section summarizes in the de Sitter (dS) case the basic semiclassical relations and resulting leading IR effects on propagators, due to tensor loops. Section three then summarizes checks of these results, through explicit loop calculations of in-in correlators, in which the leading IR behavior is extracted.  Section four extends the semiclassical relations to slow-roll spacetimes, accounting for both tensor and scalar loops, and performs another check on the formalism based on results in the literature.  There is also a brief comment on the relation to the $\\delta N$ approach.  Section five calculates the leading corrections to cosmological observables, such as the tensor to scalar ratio, and nongaussianity parameters.  It also contains a worked example of scenarios of slow-roll following chaotic inflation, where such corrections become large due to long-time evolution.  Section six closes with further discussion of interpretation of these IR effects, and the parallels with black hole physics.  The appendices contain review of previously known consistency relations, more complete derivation of the semiclassical consistency relations that we use, and more details of the in-in loop calculations. ",
        "conclusions": ""
    },
    "1005/1005.3323_arXiv.txt": {
        "abstract": "We have searched $Microvariability ~ and ~ Oscillations ~of ~STars$ (MOST) satellite photometry obtained in 2004, 2005, and 2007 of the solar-type star HD 209458 for Trojan asteroid swarms dynamically coupled with the system's transiting ``hot Jupiter'' HD 209458b.  Observations of the presence and nature of asteroids around other stars would provide unique constraints on migration models of exoplanetary systems. Our results set an upper limit on the optical depth of Trojans in the HD 209458 system that can be used to guide current and future searches of similar systems by upcoming missions. Using cross-correlation methods with artificial signals implanted in the data, we find that our detection limit corresponds to a relative Trojan transit depth of 1$\\times10^{-4}$, equivalent to $\\sim$1 lunar mass of asteroids, assuming power-law Trojan size distributions similar to Jupiter's Trojans in our solar system. We confirm with dynamical interpretations that some asteroids could have migrated inward with the planet to its current orbit at 0.045 AU, and that the Yarkovsky effect is ineffective at eliminating objects of $>$ 1 m in size. However, using numerical models of collisional evolution we find that, due to high relative speeds in this confined Trojan environment, collisions destroy the vast majority of the asteroids in $<$10 Myr. Our modeling indicates that the best candidates to search for exoTrojan swarms in 1:1 mean resonance orbits with ``hot Jupiters\" are young systems (ages of about 1 Myr or less). Years of Kepler satellite monitoring of such a system could detect an asteroid swarm with a predicted transit depth of $3 \\times 10^{-7}$. ",
        "introduction": "\\indent Currently, there are slightly over 440 known exoplanets (The Extrasolar Planets Encyclopaedia; exoplanet.edu). About 60 of these have orbital planes aligned with our line of sight so that the planets transit their parent stars causing small dips in brightness repeating at the planet's orbital period. Most of the known exoplanetary systems have gas giant planets in small, very short-period orbits, dubbed ``hot Jupiters\" due to their large sizes and proximities to their parent stars.  Dynamical models \\citep{laughcham2002, thommes2005, cressnel2006, lyra2009} predict that some of these hot Jupiters may be accompanied in their orbits by swarms of Trojan asteroids at the L4 and L5 Lagrangian points, like those in resonance with Jupiter in our own solar system \\citep{stacon2007}. The detection and characterization of Trojan asteroids in an exoplanetary system would provide important constraints on the dynamical evolution of giant exoplanets.  \\indent Only a few years ago, the prospect of being sensitive to even a lunar mass of asteroids around another star would have seemed like science fiction.  The new detection limits made possible by space-based photometry missions like MOST, CoRoT, and Kepler mean for the first time that planetary scientists can seriously explore limits on the formation and evolution of possible exoTrojan swarms.  \\indent One way of detecting exoplanetary Trojans is through transit timing by combining radial velocity observations and photometric observations of a transiting planet. Through this approach a Trojan swarm can be detected via a time lag between the radial velocity null and the time of the central transit. Using this method, \\citet{forgaudi2006} ruled out Trojan companions to HD 209458b of total mass greater than $\\sim$ 13 M$_{\\earth}$ at a 99.9\\% confidence level.  \\indent In this paper, we present a more direct and more sensitive approach, by searching directly for the transit signal of Trojan swarms with precise spacebased photometry. Our target for this study was the HD 209458 system, which contains a hot Jupiter (designated HD 209458b). The planet was discovered in radial velocity measurements by \\citet{mazeh2000} and the transits were first reported by \\citet{charbonneau2000}.  \\indent HD209458b has a circular orbit of semimajor axis $a = 0.045$ AU and period $P = 3.52474859 \\pm 0.00000038$ days \\citep{knutson2007}. The primary HD 209458 is a relatively-bright ($V = 7.65$) G0 star with $T_{eff} = 6000 \\pm 50$ K, luminosity $L = 1.61$ L$_{\\odot}$ \\citep{mazeh2000}, mass $M = 1.10 \\pm$ 0.07 M$_{\\odot}$, and radius $R = 1.13 \\pm$ 0.02 R$_{\\odot}$ \\citep{knutson2007}. Although quite uncertain, \\citet{melo2006} give an estimate of 3 Gyr for the age of the system.  HD 209458 is a good starting point for the steps to compare the dynamical state of asteroids in an exoplanetary system with the current state of our solar system: the host star is Sun-like, and it is one of the brightest transiting systems available for study.  \\indent It is widely accepted that hot Jupiters formed at orbital distances greater than their current locations and then migrated inward. We explore Trojan survival in such a system. To be visible today, the Trojan populations would have first had to survive a migration from about 5 AU down to about 0.05 AU. Once they arrived at this distance, both self-collisions and radiation effects (such as the Yarkovsky effect, which can cause small objects to undergo orbital changes) may be important for these hypothetical ``hot Trojans'' over times scales comparable to the system's age. We explore these effects through numerical simulations and discuss the results.  \\indent This paper is organized as follows. In Section 2 we describe how the photometric data were collected. In Section 3, we describe the data analysis, methods, and results. Section 4 is an estimate of the photometric cross section of the Trojan clouds in the HD 209458 system based on the solar system Trojan size distribution. Section 5 presents numerical calculations regarding dynamical and collisional evolution, as well as the Yarkovsky effect in this system. Finally, Section 6 provides conclusions based on these results. ",
        "conclusions": "\\indent In this paper, photometric data on HD 209458 from the MOST space telescope were analyzed with the intent of finding Trojan asteroids swarms in 1:1 mean motion resonance with the known planet HD 209458b. Using cross-correlation techniques on data with synthetic Trojan cloud transits, we were able to set an upper limit on the Trojan transit depth of 1$\\times10^{-4}$, which corresponds to an upper limit in asteroid swarm mass to $\\sim$ 1 lunar mass.  We then assessed the dynamical effects on such a Trojan cloud in the HD 209458 system. The main findings are as follows.  \\begin{itemize}  \\item During the migration of the Trojan swarm (as they migrate with the planet from $\\sim$ 5 AU to 0.045 AU), the libration amplitudes grow by a factor of $\\sim$ 3.3 leaving $\\sim 10\\%$ of the Trojans surviving migration.  \\item The collisional evolution of the cloud grinds it down to below 100 m size objects in $\\ll$10 Myr and reduces its surface area 10,000 fold. Neither increasing nor decreasing the initial mass of the cloud changes the results significantly. After only a few Myr the total cross section remaining in the Trojan swarm (in bodies larger than tens of meters) has dropped to $<10^{-8}$ that of the star.  Based on the initial photometric performance of Kepler long cadence data \\citep{jenkins2010}, for a solar-type star of V $\\sim$ 10 with Trojan asteroid swarms in a 3.5 day orbit, the light curve would reach a sensitivity to transit depth of $3 \\times 10^{-7}$.  In a young system, this would be sufficient to detect the exoTrojan swarms based on our model predictions.  \\item Radiation forces do not cause a semimajor axis drift of the leftover small objects, only a shift in the location of their libration centers. Further collisions between these small objects could then cause a random walk in their libration amplitudes, eventually destabilizing the Trojans and eliminating them on a time scale of 2 Myr.  \\item A nearly opaque cloud might result from the collisionally ground asteroid population just before the radiation forces take effect. Such a cloud would have a very strong photometric signature for a brief time period ($<$ 1 Myr) due to the dense population of small particles. This would occur during the collisional phase of the cloud evolution, and since this is a fast process, this could only be observed in a very young system ($\\ll$100 Myr).  \\end{itemize}  \\indent Thus, the best chance of detecting Trojan asteroid swarms in another solar system in the foreseeable future would be in very young systems, where either the collisional depletion has not yet reduced the surface area significantly, or where the cloud turned opaque for a short time due to the high density of centimeter-sized objects before elimination by Poynting-Robertson drag. Detection in an older system would require a recent breakup of a large body that generated a collisional cascade of smaller objects.  \\indent Our simulations do rely on the assumption of simple migration of the gas giant planet.  Other processes, involving planet-planet interactions and later tidal circularization of the exoplanet's small orbit, could lead to different scenarios of exoTrojan migration and evolution.  However, we note that most of the observations of the Rossiter-McLaughlin effect in exoplanet systems, including HD 209458 \\citep{winn2005}, have pointed to only small levels of spin-orbit misalignment in systems with close-in giant planets \\citep{winn2007, narita2009}.  There is no reason to strongly suspect planet-planet interactions in the case of HD 209458 and other hot Jupiter systems with small, nearly circular orbits.  \\indent What about exoplanets in orbits with much larger semi-major axes?  For exo-Jupiters with relatively large semi-major axes, the prospects of detecting transits in the existing sample are limited. The Kepler mission will expand the sample and should change that situation in the coming years. Many of the known exoplanets with larger semi-major axes also have high eccentricities. In these systems, there could have been dynamical interactions with other planets earlier in the histories of the systems.  We argue that these are not systems where it would be fruitful to search for Trojans in resonant orbits, and for that reason, we did not explore this broader range of parameter space in the numerical simulations presented in this paper.  \\indent Our analysis of ultra-precise MOST photometry provides the most sensitive upper limit yet published on the amount of Trojan asteroidal material in an exoplanetary system. This upper limit inspired us to perform numerical simulations to show what Trojan optical depth might be present.  By estimating the expected optical depth of an exoplanetary asteroid cloud, we have set meaningful limits on future photometric searches for exoTrojan swarms by missions such as Kepler, and focus attention in particular on very young systems."
    },
    "1005/1005.3115_arXiv.txt": {
        "abstract": "{  Our PACS and SPIRE images of the Aquila Rift and part of the Polaris Flare regions, taken during the science demonstration phase of \\emph{Herschel} discovered fascinating, omnipresent filamentary structures that appear to be physically related to compact cores. We briefly describe a new multi-scale, multi-wavelength source extraction method used to detect objects and measure their parameters in our \\emph{Herschel} images. All of the extracted starless cores (541 in Aquila and 302 in Polaris) appear to form in the long and very narrow filaments. With its combination of the far-IR resolution and sensitivity, \\emph{Herschel} \\emph{directly} reveals the filaments in which the dense cores are embedded; the filaments are resolved and have deconvolved widths of $\\sim$ 35{\\arcsec} in Aquila and $\\sim$ 59{\\arcsec} in Polaris ($\\sim$ 9000\\,AU in both regions). Our first results of observations with \\emph{Herschel} enable us to suggest that in general dense cores may originate in a process of fragmentation of complex networks of long, thin filaments, likely formed as a result of an interplay between gravity, interstellar turbulence, and magnetic fields. To unravel the roles of the processes, one has to obtain additional kinematic and polarization information; these follow-up observations are planned.  } ",
        "introduction": "\\label{Introduction}  \\emph{Herschel} \\citep{Pilbratt_etal2010} provides the best opportunity to study the earliest stages of star formation and the origin of the stellar initial mass function (IMF). Prestellar cores and young (Class 0) protostars emit the bulk of their luminosities at wavelengths ~80--400\\,{\\um}, which makes the \\emph{Herschel} imaging instruments PACS \\citep{Poglitsch_etal2010} and SPIRE \\citep{Griffin_etal2010} perfect for performing a census of these objects down to ~0.01--0.1\\,{\\Msun} in the nearby (distances $D \\la$ 500 pc) molecular cloud complexes. Our \\emph{Herschel} key project \\citep[Gould Belt survey,][]{Andre_etal2010} aims at probing the link between diffuse cirrus-like structures and compact cores with the main scientific goal to elucidate the physical mechanisms of the formation of prestellar cores out of the diffuse medium, which is crucial for understanding the origin of stellar masses.  During the science demonstration phase of \\emph{Herschel} observations we imaged two large fields, the Aquila Rift and part of the Polaris Flare (hereafter called Aquila and Polaris fields for brevity), with the aim of better understanding star formation in extremely different environments. The high-quality far-infrared images are extremely rich in new information on the earliest phases of star formation, which will be explored in full detail in the future. In this short paper we focus on merely one discovery made by \\emph{Herschel}: the impressive, ubiquitous network of filamentary structures in the interstellar medium (ISM) and its remarkably close physical relationship with the objects that appear to form within the filaments. ",
        "conclusions": "\\label{results.and.discussion}  To derive properties of the compact objects in the Aquila and Polaris fields, we performed source extraction with \\textit{getsources} (Sect.~\\ref{source.extraction.method}) and made a careful analysis of the resulting extraction catalogs, selecting most reliable objects and determining core mass function for both fields \\citep{Ko\"nyves_etal2010,Andre_etal2010}. To assess completeness levels of our observations, we created the synthetic skies with the actual observed background of our images (after removing all extracted objects) and populated them with cores and protostars from spherical radiative transfer models \\citep{Men'shchikov_etal2010}, then performed extractions on the skies.  The filamentary structures are clearly visible in the left panels of Fig.~\\ref{figure1}, where both the Aquila and Polaris fields are shown as 3-color composites created from the images at 500, 160, and 70\\,{\\um}. With red color, the images highlight strong density enhancements and low temperatures, whereas the material within two bright hot regions (HII region W\\,40 in the north and around MWC\\,297 and Sh\\,2-62 stellar group in the south) of the Aquila field is visible in green, blue, or white. From this one could immediately conclude where to look for prestellar cores, and it is indeed the red filaments where the reddest cores are visible. The colors are not very pronounced in the Polaris field, suggesting much smaller temperature and density variations across the field, even in its densest filaments. To better visualize and characterize the filaments, we utilized several approaches.  High-contrast images of the cores and filaments in both fields were obtained as part of our source extraction technique by applying the successive unsharp masking (Eq.~\\ref{successive.unsharp.masking}) with $f_{\\rm s}$=1.05, $G_1$=9{\\farcs}39, and $N$=99. In particular, the right panels in Fig.~\\ref{figure1} show that the visibility of the very narrow and long filaments is dramatically improved in some of our ``single-scale'' images (which effectively select a narrow range of scales, with both larger and smaller scales filtered out). For the sake of illustration, a 40{\\arcsec} scale (in which the structures are best visible) is displayed in Fig.~\\ref{figure1}. The compact objects are populating the filaments almost like pearls on threads in a necklace, giving them a nodular, lumpy appearance. The same decomposition at all \\emph{Herschel} wavelengths proved that the same structures are seen across the entire range, from 70 to 500\\,{\\um}. There are, however, noticeable changes in the appearance of the filaments in Aquila (around W\\,40) at the shortest wavelengths (Fig.~\\ref{figure2}, upper left), presumably due to increasing opacity and to changes in temperature of some of the filaments (complex three-dimensional structures).  We also extracted filamentary structure from the observed images with the MCA software \\citep[morphological component analysis,][]{Starck_etal2004}. The idea of the method is to morphologically decompose a signal into its building blocks, which can be represented by isotropic wavelets, ridgelets, or curvelets. Filamentary structures were separated from the more isotropic features by applying both the curvelet and wavelet transforms in 100 iterations (with exponentially decreasing thresholds) to achieve convergence in both the wavelet and curvelet representations of the original images. The curvelet component \\citep[Fig.~\\ref{figure2}, right panels; see also][]{Andre_etal2010} shows filaments very similar to those displayed in Fig.~\\ref{figure1} and left panels of Fig.~\\ref{figure2}; we also get analogous results when processing the column density images \\citep[the latter are shown in][]{Ko\"nyves_etal2010}.    Measurements of the thickness of several high-contrast filaments in our column density maps (obtained with a uniform 39{\\arcsec} resolution of the 500\\,{\\um} band) give 48{\\arcsec} for Aquila and 66{\\arcsec} for Polaris (FWHM), uncertain to within 12{\\arcsec}; when deconvolved, they are 28{\\arcsec} and 53{\\arcsec}. If we adopt the distances to the regions of 260\\,pc \\citep{Straizys_etal2003,Bontemps_etal2010} and 150\\,pc \\citep{HeithausenThaddeus1990}, the deconvolved linear widths are $\\sim$ 7300\\,AU and 8000\\,AU. Measurements of the widths at several locations across a few well-behaved filaments (those not blended with the other filaments) in the observed images resulted in the values of 30, 36, 42, 42, 54{\\arcsec} (for Aquila at the PACS+SPIRE wavelengths) and 63, 64, 69{\\arcsec} (for Polaris at the SPIRE wavelengths). The filaments are more resolved at PACS wavelengths and the values would imply average deconvolved widths of 35{\\arcsec} for Aquila and 59{\\arcsec} for Polaris (9000\\,AU in both fields). The largest FWHM sizes of the extracted cores are comparable to the filaments' widths (within $\\sim$50{\\%}) for both fields, confirming the visual impression that the objects are largely confined to the filaments they are embedded in. There seems to be a wide distribution of the lengths of the filaments; some of them are as long as $\\sim$0{\\fdg}5 (Figs.~\\ref{figure1}, \\ref{figure2}), which would correspond to a few pc. It is clear, however, that one needs to perform a more systematic and complete study of the filaments' properties; that will be a subject of a future work.  The Aquila and Polaris fields present very different physical conditions; the former shows very active star formation, while the latter displays almost none. Indeed, we find many extremely cold and dense filaments ($T_{\\rm{d}}\\approx$ 7.5--15\\,K, $N_{\\rm{H_2}}\\approx 5\\times 10^{20}$--$1.4\\times 10^{23}$\\,cm$^{-2}$) in Aquila; most of them are gravitationally unstable and fragmented into several hundreds of prestellar cores \\citep{Andre_etal2010}. In contrast, the Polaris field contains warmer filaments with much lower column densities ($T_{\\rm{d}}\\approx$ 10--15\\,K, $N_{\\rm{H_2}}\\approx 3\\times 10^{20}$--$8.6\\times 10^{21}$\\,cm$^{-2}$) and we find no clear examples of prestellar cores \\citep{Andre_etal2010,Ward-Thompson_etal2010}. To illustrate the very close relationship between the filaments and detected objects, we overplot locations of the latter on the curvelet transform for both fields. An inspection of Fig.~\\ref{figure2} shows that all objects detected in Aquila and Polaris are indeed located within very narrow filaments. Our maps of a stability parameter of the filaments (their mass per unit length) suggest that the Polaris filaments are stable, while many dense filaments in Aquila are gravitationally unstable \\citep{Andre_etal2010}.  Several possible models for the formation of filamentary cloud structures have been proposed in the literature. In particular, numerical simulations of supersonic magnetohydrodynamic turbulence in weakly magnetized clouds always generate complex systems of shocks, which fragment the gas into high-density sheets, filaments, and cores \\cite[e.g.][]{Padoan_etal2001}. Filamentary structures are also produced in turbulent simulations of more strongly magnetized molecular clouds, including the effects of ambipolar diffusion \\citep{NakamuraLi2008}. Another possibility is that the filaments originate from gravitational fragmentation of sheet-like cloud layers \\citep{Nagai_etal1998}. Interestingly, filaments can result from purely gravitational amplification of a random field of initial density fluctuations as is the case in large cosmological simulations \\citep[e.g.][]{Ocvirk_etal2008}. Our measurements of the radial profiles of selected filaments in Aquila and Polaris suggest that magnetic fields may play an important role in shaping the structure of the observed filaments, if not in forming the filaments. Indeed, they indicate that the density falls off approximately as $\\rho\\propto r^{-1.5}$ (in Aquila) or $\\rho\\propto r^{-2}$ (in Polaris) away from the axes of the filaments. In both cases, the measured radial density profiles are much shallower than the steep $\\rho\\propto r^{-4}$ profile of unmagnetized hydrostatic filaments \\citep{Ostriker1964}, as already pointed out by \\citet{Lada_etal1999} and \\citet{JohnstoneBally1999} for two specific filaments. Models of unmagnetized equilibrium filaments are therefore inconsistent with observations. By contrast, models of equilibrium filaments with primarily toroidal or helical magnetic fields can account for radial profiles ranging from $\\rho\\propto r^{-1}$ to $\\rho\\propto r^{-2}$, in agreement with observations \\citep{FiegePudritz2000}.  Filaments are seen in numerous astronomical images and the filamentary structure of molecular clouds has been known for some time \\citep[e.g.][]{SchneiderElmegreen_1979,Goldsmith_etal2008}. However, \\emph{Herschel} with its combination of the far-IR resolution and sensitivity \\emph{directly} reveals the filaments in which the cores are forming. Our first observations demonstrate the fascinating, omnipresent filamentary density structure of the cold ISM and its extraordinarily intimate physical relationship with the objects that form within the filaments. They enable us to suggest that in general dense cores may originate in a process of fragmentation of long, thin filaments, presumably formed as a result of an interplay between gravity, interstellar turbulence, and magnetic fields \\citep[see also discussion in][]{Andre_etal2010}. To unravel the role and relative importance of these processes, one needs additional kinematic and polarization information \\citep[cf.][]{Matthews_etal2001}; these follow-up observations are planned."
    },
    "1005/1005.1921_arXiv.txt": {
        "abstract": "{\\PSpectRe\\ is a C++ program that uses Fourier-space pseudo-spectral methods to evolve interacting scalar fields in an expanding universe.   \\PSpectRe\\ is optimized for the analysis of parametric resonance in the post-inflationary universe and provides an alternative to finite differencing codes, such as {\\sc Defrost} and {\\sc LatticeEasy}.  \\PSpectRe\\ has both second-  (Velocity-Verlet) and fourth-order (Runge-Kutta) time integrators. Given the same number of  spatial points and/or momentum modes,  \\PSpectRe\\ is not significantly slower than finite differencing codes, despite the need for multiple Fourier transforms at each timestep, and exhibits excellent energy conservation.  Further, by computing the post-resonance equation of state, we show that in some circumstances  \\PSpectRe\\  obtains reliable results while using  substantially fewer points than a finite differencing code.     \\PSpectRe\\ is designed to be easily extended to other problems in early-universe cosmology, including the generation of gravitational waves during phase transitions and pre-inflationary bubble collisions.  Specific applications of this code will be described in future work. }   \\begin{document} ",
        "introduction": "A key challenge for any inflationary model  is to ensure that inflation {\\em ends\\/} and that the post-inflationary universe thermalizes, setting the stage for a conventional, hot-big-bang cosmology. A promising approach to this problem is provided by preheating and parametric resonance, in which the post-inflationary universe undergoes explosive particle production \\cite{Traschen:1990sw,Shtanov:1994ce,Kofman:1995fi}.  Preheating  leads to a fascinating range of phenomenological consequences, possibly including superheavy dark matter \\cite{Chung:1998zb}, gravitational-wave production \\cite{Khlebnikov:1997di,Easther:2006gt,Easther:2006vd,Easther:2007vj,GarciaBellido:2007dg,GarciaBellido:2007af,Dufaux:2007pt} and non-Gaussianity~\\cite{Chambers:2007se,Chambers:2008gu,Bond:2009xx}. Resonance can occur in a broad swathe of cosmological models and, since it depends on the coupling between the inflaton sector and the ``rest of physics,'' any observable consequences of  resonance constitute a unique window into the inflationary epoch.  Given that resonance is an intrinsically nonlinear process, understanding it is a challenging task, and frequently requires numerical studies. Consequently, expanding on a theoretical framework initially developed by Khlebnikov and Tkachev~\\cite{Khlebnikov:1996mc}, a number of codes have been developed to study scalar field dynamics in an expanding universe, including \\textsc{LatticeEasy} \\cite{Felder:2000hq} (not to mention the parallelized \\textsc{ClusterEasy} \\cite{Felder:2007nz}, and GPU accelerated \\textsc{CUDAEASY} \\cite{Sainio:2009hm}) and \\textsc{Defrost}\\cite{Frolov:2008hy}. These programs all use  finite differencing methods to compute spatial derivatives of the scalar fields. In this paper we describe \\PSpectRe, a Fourier-space spectral code for solving the Klein-Gordon equations of two scalar fields in an expanding universe with an arbitrary quartic interaction potential. Technically, this is a pseudo-spectral code, in that nonlinear terms in the potential and its derivatives are computed by first converting the fields into their position-space representation, performing the necessary multiplications, and then taking the inverse transform \\cite{BoydBK}. Since spatial derivatives are taken by multiplying the corresponding Fourier component by the wavenumber $k_i$, pseudo-spectral methods are free of differencing noise\\footnote{By differencing noise we mean the error introduced when computing approximations to a derivative by subtracting neighboring-point values and then dividing by some multiple of the grid spacing. Since neighboring points often have similar values, subtracting them using finite-precision, floating-point arithmetic often yields a result with few significant digits. When the result is divided by some multiple of the grid spacing, the rounding error is greatly magnified, resulting in a significant source of numerical error.}, and have excellent spatial resolution~\\cite{BoydBK}. On the other hand, overhead from the Fourier transforms and inverse transforms exacts a cost in terms of computational efficiency.  As currently implemented, \\PSpectRe\\ allows both second- and fourth-order-accurate time integration steps. The latter is more computationally intensive, but  delivers excellent energy conservation.    \\PSpectRe\\ is currently implemented to run on a single, shared memory machine, but threads efficiently across 16 (effective) cores.  The rate limiting step   is the fast Fourier transform  (FFT) algorithm. It would be straightforward to port \\PSpectRe\\ to a cluster environment, given that the FFT libraries we use exist in a fully parallelized form.  The goal of this paper is to introduce and describe \\PSpectRe, and to demonstrate that it provides a useful cross check on existing finite-differencing-based codes. Looking to the future, spectral codes have excellent accuracy and precision, and the investigation of preheating-generated observables such as gravitational waves and non-Gaussianity require accurate solutions of the background dynamics. The initial version of  \\PSpectRe\\ is equivalent to \\textsc{LatticeEasy} and \\textsc{Defrost} in that it solves the Klein-Gordon equation in a rigid, expanding background. However,  Easther, Giblin and Lim have previously used a spectral algorithm to compute the gravitational-wave background generated during resonance  \\cite{Easther:2006vd,Easther:2007vj}.  This approach  naturally lends itself  to implementation within \\PSpectRe, and will be pursued in future work.  Finally, in earlier work Easther and Parry used a pseudo-spectral code to study gravitational effects on resonance in a 1+1 dimensional scenario \\cite{Parry:1998pn,Easther:1999ws}. In terms of performance, a ``stock'' version of {\\sc Defrost}  is a few times faster than \\PSpectRe\\ on the same hardware (with a $256^3$ grid), but the difference is not dramatic.    However, we will compute the post-inflationary equation of state, and show that \\PSpectRe\\ can achieve reliable results with a $64^3$ grid in circumstances where {\\sc Defrost\\/} requires $128^3$ or even $256^3$ points.  We attribute this improvement to the lack of differencing noise in spectral codes.   The format of this paper is as follows. We first review the scalar field evolution equations, and the specific program variables used within \\PSpectRe, which are the same as those employed by  \\textsc{LatticeEasy}. We then describe \\PSpectRe's basic computational engine, and present results on energy conservation and the spatial distribution of energy following resonance for the well-studied scenario of an inflation, $\\phi$, with a quadratic potential, coupled to a massless field, $\\chi$, via a $\\phi^2 \\chi^2$ interaction term. ",
        "conclusions": "We have demonstrated that \\PSpectRe\\ yields a highly-accurate and computationally-efficient approach to the solution of scalar-field evolution during parametric resonance.  \\PSpectRe\\ thus provides a useful contrast to  finite-differencing  algorithms  such as those in {\\sc Defrost} and {\\sc LatticeEasy}.  The present discussion is intended primarily as a ``proof of concept'', demonstrating  that  a pseudo-spectral algorithm is a practical method for the analysis of preheating, despite the extra computational overhead introduced by the nonlinear terms in the evolution equations.  In doing so, we have recovered a variety of existing results in the literature.     We find that \\PSpectRe's energy conservation performance is excellent, and scales in well-understood ways as the grid size and timestep are varied.  Moreover, we demonstrate that some energy is ``lost'' to the simulation when it is transferred to high frequency modes. While  \\PSpectRe\\ does not evolve these modes, it can account for them by ``padding'' the spatial integral used to compute the potential energy.  As a pseudo-spectral code,  \\PSpectRe\\ does not rely on comparing values of fields at adjacent lattice points to take spatial derivatives. Consequently, in some circumstances  \\PSpectRe\\ provides accurate results with smaller grids than those needed by finite differencing codes.   When the grid sizes are equal, \\PSpectRe\\ is still competitive with finite-differencing codes in terms of runtime, but the repeated Fourier transforms do introduce a noticeable performance penalty.  Reducing the grid size by  a factor of $2$ naively improves runtime by (slightly more than) a factor of $8$, with fixed timestep, while maintaining the same effective resolution as a finite-differencing code.  Consequently, by changing runtime settings, \\PSpectRe\\ can be optimized for precision,  speed, or the ability to run simulations over large spatial volumes.   \\appendix"
    },
    "1005/1005.3924_arXiv.txt": {
        "abstract": "{The Rosette molecular cloud is  promoted as the archetype of a triggered star-formation site. This is mainly due to its morphology, because the central OB cluster NGC~2244 has blown a circular-shaped cavity into the cloud and the expanding \\hii-region now interacts with the cloud. } {Studying the spatial distribution of the different evolutionary states of all star-forming sites in Rosette and investigating possible gradients of the dust temperature will help to test the 'triggered star-formation' scenario in Rosette.} {We use continuum data obtained with the PACS (70 and 160 $\\mu$m) and SPIRE instruments (250, 350, 500 $\\mu$m) of the \\emph{Herschel} telescope during the Science Demonstration Phase of HOBYS.} {Three-color images of Rosette impressively show how the molecular gas is heated by the radiative impact of the NGC~2244 cluster. A clear negative temperature gradient and a positive density gradient (running from the \\hii-region/molecular cloud interface into the cloud) are detected. Studying the spatial distribution of the most massive dense cores (size scale 0.05 to 0.3 pc), we find an age-sequence (from more evolved to younger) with increasing distance to the cluster NGC~2244. No clear gradient is found for the clump (size-scale up to 1 pc) distribution.} {The existence of temperature and density gradients and the observed age-sequence imply that star formation in Rosette may indeed be influenced by the radiative impact of the central NGC~2244 cluster. A more complete overview of the prestellar and protostellar population in Rosette is required to obtain a firmer result. } ",
        "introduction": "\\label{intro} The Rosette is a spectacular region due to its well-known emission nebula, powered by the central OB cluster NGC~2244 (Figs.~\\ref{rosette-3col}, \\ref{overview}) that has blown a cavity into the molecular cloud. The expanding \\hii-region is interacting with a high-mass (few times $10^5~M_\\odot$) star-forming molecular cloud at 1.6~kpc from the Sun (Williams et al. \\cite{williams1994}, Schneider et al. \\cite{schneider1998a}, Heyer et al. \\cite{heyer2006}), creating photon dominated regions (PDRs) along the interface (Schneider et al. \\cite{schneider1998b}).  Bright PAH emission, first seen in the IRAS bands by Cox et al. (\\cite{cox1990}), indicates the presence of warm dust. Near- and mid-IR surveys (Phelps \\& Lada \\cite{phelps1997}, Li \\& Smith \\cite{li2005}, Balog et al. \\cite{balog2007}, Roman-Zuniga et al. \\cite{roman2008}, Poulton et al. \\cite{poulton2008}) investigated the IR-population of Rosette in greater detail .   Ionizing radiation has an effect on the evolution of a molecular cloud and is thus closely linked with star formation. Photoevaporation of gas on an existing (primordial) clumpy cloud structure causes the formation of {\\sl globules}. Well-known features in the Rosette molecular cloud (hereafter RMC) are the 'elephant trunks', 'speck'- and 'teardrop' globules (Schneps et al. \\cite{schneps1980}, White et al. \\cite{white1997}), which show bright rim emission. These globules can form low-mass stars and shrink slowly (e.g. Lefloch \\& Lazareff \\cite{lefloch1994}). Rayleigh-Taylor instabilities can explain the elephant trunk globules in Rosette (e.g. Patel et al. \\cite{patel1993}). Schneider et al. (\\cite{schneider1998a}) showed that basically all CO clumps close to the \\hii-region collapse and form stars. Recently, Gritschneder et al. (\\cite{grit2009}) ran hydrodynamic models including turbulence and ionizing radiation. They showed that it is due to the heating and consequently the increasing pressure of the lower-density interclump gas that primordial, denser clumps are even more compressed,  forming pillar-like structures that point towards the source of radiation. In the high-density (n$\\sim$10$^7$ cm$^{-3}$) tip of the pillars, star formation takes place. However, the extent to which star formation is promoted deeper within the cloud, where the level of ionization is drastically falling, is not clear.   Although there is a clear influence of radiation on the evolution of the RMC in general, it is less clear what its specific link to star formation is. In the classical picture of {\\sl sequential triggered} star formation (Elmegreen \\& Lada \\cite{elmlada1977}), an expanding \\hii-region assembles gas in the layer compressed between the ionization front and the shock front, which then fragments into gravitationally unstable dense cores that form stars. The ionization front of this new association propagates further into the cloud and induces the next generation of star formation.     The Rosette was proposed as an archetype of a triggered star formation site. Many other studies (e.g. Deharveng et al. \\cite{deharveng2005}) have revealed possible candidates on smaller size scales with a clearly-defined, simple geometry and indications of star-formation in a swept-up shell of gas around a central \\hii-region.  The Rosette is a more demanding region, because it is a massive molecular cloud with a complex structure caused by the presence of at least three clusters that illuminate the \\hii-region (NGC~2244, NGC~2237, and the recently detected RMC-XA (Wang et al. \\cite{wang2009}). There are arguments that star-formation in Rosette may be the result of a triggering process: \\\\ (i) Balog et al. (\\cite{balog2007}) and Roman-Zuniga et al. (2008) found that the average NIR excess fraction for stellar clusters {\\sl increases} with distance to the cluster center. (ii) CO line studies of Williams et al. (\\cite{williams1994}) show that star-formation activity is more intense in the \\hii-region/molecular cloud interface region than in the molecular cloud center.   However, the triggered star-formation scenario is likely not {\\sl sequential}. Roman-Zuniga et al. (\\cite{roman2008}) point out that the relative age differences of the clusters are not consistent with a sequential triggered star formation scenario. They propose that the overall age sequence of cluster formation could be primordial, possibly resulting from the formation and evolution of the molecular cloud itself.  It is thus unclear whether the star-formation activity in Rosette is progressing from the \\hii-region interface into the cloud center, or whether it is a global process with a random spread of evolutionary stages and no obvious link between more and less evolved star-forming sites.       \\begin{figure}[ht] \\includegraphics[angle=0,width=90mm]{overview_spire_paper.jpg} \\caption [] {\\emph{Herschel}/SPIRE map at 250 $\\mu$m of Rosette. Individual regions referred to in the text are numbered and labeled. The stars indicate the seven O-stars (O4 to O9) of NGC~2244.} \\label{overview} \\end{figure}  \\section {Observations}  \\label{obs}  The RMC was chosen as a representative source (in terms of mass, star-formation activity etc.) for the HOBYS (\\emph{Herschel} imaging survey of OB Young Stellar objects) guaranteed time key program (Motte et al. 2010) on the \\emph{Herschel} satellite (Pilbratt et al. 2010), using the SPIRE (Griffin et al. 2010, Swinyard et al. (\\cite{swinyard2010}), and PACS (Poglitch et al. 2010) continuum arrays.  The SPIRE and PACS data from 70 to 500 $\\mu$m were obtained on october 20th 2009 in the parallel mode with a scanning speed of 20$''$/sec. Two cross-linked coverages of size 1$^\\circ$45$'\\times$1$^\\circ$25$'$ were performed, focussing on the southeast corner of the Rosette nebula, i.e. the main part of the molecular cloud (see Fig.~\\ref{overview}) The SPIRE data were reduced with HIPE 2.0, using modified versions of the pipeline scripts (for example observations that were taken during the turnaround at the map borders were included). A median baseline and the `naive-mapper' (i.e. a simple averaging algorithm) were applied to the data.        \\begin{figure}[ht] \\includegraphics[angle=-90,width=90mm]{rosette_temp.jpg} \\caption [] {Dust-temperature map of the Rosette molecular cloud, obtained from simultaneously fitting the five observing wavebands of PACS and SPIRE. The temperature scaling in Kelvin is given on the right. Note that there is probably a zero-point offset in the maps due to the median filter on the scan legs.  Only areas that are covered by SPIRE {\\sl and} PACS are shown. Black crosses and label 'PL' indicate the position of the clusters identified by Phelps \\& Lada (\\cite{phelps1997}), red crosses and label 'A','B' etc. the clusters from Poulton et al. (\\cite{poulton2008}), and label 'REFL' the clusters from Roman-Zuniga et al. (2008)}. \\label{temp} \\end{figure} ",
        "conclusions": "\\label{discussion}  The prestellar and protostellar population of Rosette observed with \\emph{Herschel} was identified with two source extraction techniques that are both based on the principle that emission on large spatial scales is filtered out in order to focus on compact sources. Motte et al. (2010) discuss the core population (prestellar, protostellar, and warm cores on a size scale 0.05--0.3 pc). Di~Francesco et al. (2010) present the {\\sl clump} distribution (up to 1 pc), which is either starless or associated with YSO.    Figure~1 in Di~Francesco et al. (2010) shows the spatial distribution of starless and protostellar {\\bf clumps}. The latter are concentrated in the densest cloud regions, mainly close to the known star clusters, while the starless clumps are more evenly distributed.  {\\sl No} correlation between clump type and distance to the OB cluster is observed that may suggest an age-effect.   In Fig.~5 in Motte et al. (2010) the spatial distribution of the most massive dense {\\bf cores} is shown and in our Table~1 the statistics of these sources with regard to the different regions is summarized.  We adopt the classification for prestellar cores to evolved YSO defined in Motte et al. (2010) and compute the relative numbers of young vs. evolved sources in each region (last column in Table~1). The resulting sequence {\\sl may} indicate that there is an age-gradient across the RMC with the PL7 cluster being the youngest, in the remote part of the cloud, and the Monoceros Ridge and Shell Region, close to the \\hii-region, being the most evolved. An independent confirmation of the youth of cluster PL7 is given in Hennemann et al. (2010) in their protostellar envelope mass vs. bolometric luminosity diagram. Note that Roman-Zuniga et al. (2008) also detected an increasing fraction of young stars with increasing distance to the NGC~2244 cluster center and Balog et al. (\\cite{balog2007}) found that the Class II sources are concentrated at the center of NGC~2244, while the Class I sources are located further away from the cluster. On the other hand, cluster PL3 does not fit this scenario well, because it seems to be a more evolved region, but is located in the more remote part of the cloud. A tentative explanation may be that because it is more exposed to UV radiation and thus less shielded by denser gas, (like the Center and remote region) its evolution might have been accelerated by the radiative impact. However, we emphasize that there are not enough sources analyzed so far to draw a firm conclusion. A more detailed analysis with a larger sample and significant statistics is required, which is planned for the future.    \\begin{table}[htbp] \\begin{tabular}{lcccc} \\hline \\hline & d    & $<$n$>$ & $M$      &  Young/ \\\\ Region                  & [pc] &  [10$^3$ cm$^{-03}$]  &  [M$_\\odot$] &  evolved \\\\ \\hline &      &         &        &    \\\\ 7 PL7                   &  35  &  3.9    &  2080  &  3/0 \\\\ % 6 PL3                   &  27  &  1.5    &  1650  &  1/3 \\\\ % 5 AFGL961               &  25  &  9.5    &   792  &  4/1 \\\\ % 3 Center                &  20  &  0.6    & 22250  &  9/4 \\\\ % 1+2  Ext. + Mon. Ridge  &  10  &  0.6    &  4280  &  6/3 \\\\ % 4 Shell                 &  10  &  0.4    & 11780  &  1/2 \\\\ % \\end{tabular} \\caption{Age-sequence of the sources found in Rosette. The regions are ordered with decreasing distance (d) from the center of NGC~2244 (for simplicity we take as 'center' the position of cluster B in Fig.~3). The average density ($<$n$>$) and mass ($M$) were derived from the column density by determining the area of each region and assuming the same extent of the cloud along the line of sight. The last column gives the number of sources (young vs evolved).} \\label{cores} \\end{table}"
    },
    "1005/1005.4818_arXiv.txt": {
        "abstract": "We present nonlinear mean-field $\\alpha\\Omega$ dynamo simulations in spherical geometry with simplified profiles of kinematic $\\alpha$ effect and shear. We take magnetic helicity evolution into account by solving a dynamical equation for the magnetic $\\alpha$ effect. This gives a consistent description of the quenching mechanism in mean-field dynamo models. The main goal of this work is to explore the effects of this quenching mechanism in solar-like geometry, and in particular to investigate the role of magnetic helicity fluxes, specifically diffusive and Vishniac-Cho (VC) fluxes, at large magnetic Reynolds numbers ($R_{\\rm m}$). For models with  negative radial shear or positive latitudinal shear, the magnetic $\\alpha$ effect has predominantly negative (positive) sign in the northern (southern) hemisphere. In the absence of fluxes, we find that the magnetic energy follows an $R_{\\rm   m}^{-1}$ dependence, as found in previous works. This catastrophic quenching is alleviated in models with diffusive magnetic helicity fluxes resulting in magnetic fields comparable to the equipartition value even for $R_{\\rm m}=10^7$. On the other hand, models with a shear-driven Vishniac-Cho flux show an increase of the amplitude of the magnetic field with respect to models without fluxes, but only for $R_{\\rm m}<10^4$. This is mainly a consequence of assuming a vacuum outside the Sun which cannot support a significant VC flux across the boundary. However, in contrast with the diffusive flux, the VC flux modifies the distribution of the magnetic field. In addition, if an ill-determined scaling factor in the expression for the VC flux is large enough, subcritical dynamo action is possible that is driven by the action of shear and the divergence of current helicity flux. ",
        "introduction": "A crucial point in the study of astrophysical dynamos is to understand the mechanism by which they saturate. Nevertheless, a consistent description of this process has rarely been considered in mean-field dynamo (MFD) modeling and only a heuristic description is often used. An important phenomenon happens when the dynamo operates in closed or periodic domains: the turbulent contribution to the dynamo equation, i.e., the $\\alpha$ effect, decreases for large values of the magnetic Reynolds number. This process is known as catastrophic quenching and can pose a problem in explaining the generation of magnetic field in late type stars like the Sun or the Galaxy, where $R_{\\rm m}$ could be of the order of $10^9$ or $10^{15}$, respectively.  In the last few years the nature of the catastrophic quenching has been identified as a consequence of magnetic helicity conservation \\citep[for a review see][]{bs05a}. It has been found that in the nonlinear phase of the dynamo process, conservation of magnetic helicity gives rise to a magnetic $\\alpha$ effect ($\\alpha_{\\rm M}$) with a sign opposite to the inductive contribution due to the helical motions, i.e., the kinematic $\\alpha$ effect. As the production of $\\alpha_{\\rm M}$ depends on $R_{\\rm m}$, the final value of the magnetic field should also follow the same dependence. However, real astrophysical bodies are not closed systems, but they have open boundaries that may allow a flux of magnetic helicity. The shedding of magnetic helicity may mitigate the  catastrophic $\\alpha$ quenching.  These ideas have been tested in direct numerical simulations (DNS) in both local Cartesian and global spherical domains. In the former \\citep{B05,kapyla08} it has been clearly shown that open boundaries (e.g.\\ vertical field boundary conditions) lead to a faster saturation of a large-scale magnetic field compared with cases in closed domains (perfect conductor or triple-periodic boundary conditions). In the latter, it has been found that it is possible to build up large-scale magnetic fields either with forced turbulence \\citep{B05,mitra10b} or with convectively driven turbulence \\citep[e.g.,][]{brownetal10,ketal10}.  These models generally used vertical field boundary conditions.  In flux-transport dynamos \\citep{dc99,gdp08} as well as in interface dynamos of the solar cycle \\citep[e.g.][]{mgc97a,cmg97b} the quenching mechanism has been considered either through an {\\it ad hoc} algebraic equation or by phenomenological considerations \\citep{cnc04}, but most of the time the models do not consider the effects of magnetic helicity conservation. An exception is the recent paper by \\cite{cbg10}, where these effects have been considered in the context of an interface dynamo.  In general the magnetic helicity depends on time, so it is necessary to solve an additional dynamical equation for the contribution of the small-scale field to  the magnetic helicity together with the induction equation for the magnetic field. In the past few years, some effort has already been made to consider  this dynamical saturation mechanism in MFD models like in the 1D $\\alpha^2$ dynamo models presented in \\cite{bcc09}, in axisymmetric models in cylindrical geometry for the galactic $\\alpha \\Omega$ dynamo \\citep{setal06}, and also in models with spherical geometry for an $\\alpha^2$ dynamo \\citep{betal07}. The role of various kinds of magnetic helicity fluxes have been explored in several papers \\citep{bcc09,zetal06,setal06}.  Our ultimate goal is to develop a self-consistent MFD model of the solar dynamo, with observed velocity profiles and turbulent dynamo coefficients computed from the DNS. This is a task that requires intensive efforts. Hence we shall proceed step by step, starting with simple models and then including more realistic physics on the way.  In this work we will study the effects of magnetic helicity conservation in simplified $\\alpha \\Omega$ dynamo models for a considerable number of cases. More importantly, we shall perform our calculations in spherical geometry, which is appropriate for describing stellar dynamos, with suitable boundary conditions, and considering shear profiles which are a simplified version of the observed solar differential rotation. We shall also explore how magnetic helicity fluxes affect the properties of the solution. Two classes of fluxes are considered in this paper: a diffusive flux and a shear-driven or Vishniac-Cho (hereafter VC) flux \\citep{vch01}. We consider models with either radial or latitudinal shear. The effects of meridional circulation will be investigated in detail in a companion paper \\citep{cgb10}.  This paper is organized as follows: in Section~\\ref{sec.aodyn} we describe the basic mathematical formalism of the $\\alpha\\Omega$ dynamo, give the formulation of the equation for $\\alpha_{\\rm M}$ and also justify the fluxes included. In Section~\\ref{sec.model} we describe the numerical method and then, we present our results in Section~\\ref{sec.results} starting from a dynamo model with algebraic quenching to models with dynamical $\\alpha$ quenching and different fluxes. Finally, we provide a summary of this work in Section~\\ref{sec.conc}.  \\section[]{The $\\alpha \\Omega$ dynamo model} \\label{sec.aodyn}  In mean-field dynamo theory, the evolution of the magnetic field is described by the mean-field induction equation, \\begin{equation} \\label{eq1} \\frac{\\partial \\ob{\\bm{B}}}{\\partial t} = \\bm\\nabla \\times\\left(\\ob{\\bm{U}} \\times \\bm{\\ob{B}} + \\bm{\\ob{\\cal{E}}} - \\eta_{\\rm m}\\bm{\\nabla}\\times\\bm{\\ob{B}}\\right), \\end{equation} \\noindent where $\\bm{\\ob{B}}$ and $\\bm{\\ob{U}}$ represent the mean magnetic and velocity fields, respectively, $\\eta_{\\rm m}$ is the molecular diffusivity, $\\bm{\\ob{\\cal{E}}}=\\alpha \\bm{\\ob{B}} - \\etat\\mu_0\\bm{\\ob{J}}$ is the mean electromotive force obtained using a closure theory like the first order smoothing approximation, where $\\bm{\\ob{\\cal{E}}}$ gives the contribution of the small-scale components on the large-scale field, $\\alpha$ is the non-diffusive contribution of the turbulence, $\\etat$ is the turbulent magnetic diffusivity, $\\bm{\\ob{J}}=\\bm{\\nabla}\\times\\bm{\\ob{B}}/\\mu_0$ is the mean current density, and $\\mu_0$ is the vacuum permeability.  In spherical coordinates and under the assumption of axisymmetry, it is possible to split the magnetic and the velocity fields into their azimuthal and poloidal components, $\\bm{\\ob{B}}=B\\bm{\\hat{e}}_{\\phi}+\\bm{\\nabla}\\times (A\\bm{\\hat{e}}_{\\phi})$ and $\\bm{\\ob{U}}=r \\sin\\theta \\Omega \\bm{\\hat{e}}_{\\phi} + \\bm{u}_{\\rm p}$, respectively. For the sake of simplicity we shall not consider the meridional component of the flow, i.e.\\ ${\\bm u}_{\\rm p}=0$. Then, the toroidal and poloidal components of equation (\\ref{eq1}) may be written as \\begin{eqnarray} \\frac{\\partial B}{\\partial t} &\\!=\\!& \\label{eq.inb} s {\\bm B}_{\\rm p} \\cdot \\bm\\nabla \\Omega - \\left[\\bm\\nabla \\eta \\times (\\bm\\nabla \\times B \\bm{\\hat{e}}_{\\phi})\\right]_{\\phi} + \\eta D^2B\\;, \\quad, \\\\ \\frac{\\partial A}{\\partial t} &\\!=\\!& \\alpha B + \\eta D^2A\\;, \\label{eq.ina} \\end{eqnarray} where $D^2=\\nabla^2-s^{-2}$ is the diffusion operator, $\\eta=\\eta_{\\rm m}+\\etat$, $s=r \\sin \\theta$ is the distance from the axis, and $\\bm{B}_{\\rm p}=\\bm\\nabla\\times(A\\hat{e}_{\\phi})$ is the poloidal field.  The two source terms in equations (\\ref{eq.inb}) and (\\ref{eq.ina}), $s \\bm{B}_{\\rm p} \\cdot \\bm\\nabla \\Omega$ and $\\alpha B$, express the inductive effects of shear and turbulence, respectively. The relative importance of these two  effects may be quantified through the non-dimensional dynamo numbers: $C_{\\Omega}=\\Delta \\Omega L^2 / \\etat$ and $C_{\\alpha}=\\alpha_0 L / \\etat$, where $\\Delta\\Omega$ is the angular velocity different between top and bottom of the domain. Note that equations (\\ref{eq.inb}) and (\\ref{eq.ina}) are valid only in the limit $C_{\\Omega} \\gg C_{\\alpha}$, known as $\\alpha\\Omega$ dynamo.  The inductive effects of the shear may be understood as the stretching of the magnetic field lines due to the change in the angular velocity between two adjacent points. On the other hand, the kinematic $\\alpha$-effect is the consequence of helical motions of the plasma which produce screw-like motions in the rising blobs of the magnetic field. Using the first order smoothing approximation it may be expressed as: \\begin{equation} \\alpha_{\\rm K}=-\\textstyle{\\frac{1}{3}}\\tau\\overline{\\bm{\\omega}\\cdot\\bm{u}}\\;, \\end{equation} where, $\\tau$ is the correlation time of the turbulent motions and $\\bm{\\omega}=\\bm{\\nabla}\\times\\bm{u}$ is the small-scale vorticity. The saturation value of the magnetic field may be obtained by multiplying $\\alpha_{\\rm K}$ by the quenching function $f_{\\rm q}=\\left(1 + B^2/B_{\\rm eq}^2 \\right)^{-1}$, which saturates the exponential growth of the magnetic field at values close to the equipartition field strength given by $B_{\\rm eq}=(\\mu_0 \\overline{\\rho{\\bm u}^2})$. This form of algebraic quenching was introduced heuristically \\cite[see, e.g.][]{stix72} and has been often used as the standard quenching mechanism in many dynamo simulations. However, it does not give information about the back reaction process and is independent of any parameter of the system like the magnetic Reynolds number. A consistent description of the quenching mechanism will be presented in the following section.  \\subsection{Dynamical $\\alpha$ effect} \\label{s.dyn.alpha}  Recently, it has been demonstrated that when the amplitude of the magnetic field reaches values near the equipartition, the $\\alpha$-effect is modified by a magnetic contribution, the so called magnetic $\\alpha$ effect, denoted by $\\alpha_{\\rm M}$. It is usually the case that $\\alpha_{\\rm M}$ has a sign opposite to $\\alpha_{\\rm K}$ resulting thus in the saturation of the magnetic field. \\cite{pouquet76} have shown that $\\alpha_{\\rm M}$ is proportional to the small-scale current helicity of the system, hence it is possible to write $\\alpha$ as a sum of two contributions, one from the fluid turbulence and other from the magnetic field, as follows: \\begin{equation} \\alpha = \\alpha_{\\rm K} + \\alpha_{\\rm M} = -{\\textstyle\\frac{1}{3}}\\tau\\overline{\\bm{\\omega}\\cdot \\bm{u}} +\\textstyle{\\frac{1}{3}}\\tau\\overline{\\bm{j}\\cdot\\bm{b}}/\\overline{\\rho} \\;, \\label{eq.tot.alpha} \\end{equation} \\noindent where $\\overline{\\rho}$ is the mean density of the medium, assumed here as a constant, and $\\bm{j}=\\bm\\nabla\\times\\bm{b}/\\mu_0$ is the current density of the fluctuating field. The mathematical expression that describes the evolution of $\\alpha_{\\rm M}$ may be obtained by taking into account the magnetic helicity evolution \\citep{bb02}, which leads to: \\begin{equation} \\label{eq.am} \\frac{\\partial \\alpha_{\\rm M}}{\\partial t} = -2 \\etat k_{\\rm f}^2 \\left(\\frac{\\bm{\\overline{\\cal{E}}}\\cdot \\overline{{\\bm B}}}{B_{\\rm eq}^2} + \\frac{\\alpha_{\\rm M}}{\\Rm} \\right) - \\bm\\nabla \\cdot \\bm{\\overline{{\\cal{F}}}}_{\\alpha}\\;, \\end{equation} \\noindent where $k_{\\rm f}=2 \\pi / (L - r_{c})$ with $r_c=0.7L_0$ is a suitable choice for the wave number of the forcing scale, the magnetic Reynolds number $R_{\\rm M}=\\etat/\\eta_{\\rm m}$ and $\\bm{\\overline{\\cal{F}}}_{\\alpha}$ is the flux of the magnetic $\\alpha$ effect related to the flux of the small-scale magnetic helicity, $\\bm{\\overline{F}}_{\\rm f}$ through: \\begin{equation} \\label{eq.flux} \\bm{\\overline{\\cal F}}_{\\alpha} = \\frac{\\mu_0 \\overline{\\rho} \\etat k_{\\rm f}^2}{B_{\\rm eq}^2}\\bm{\\overline{F}}_{\\rm f} \\quad , \\end{equation} According to previous authors $\\alpha_{\\rm M}$ has a finite value in the interior of the domain in absence of fluxes ($\\bm{\\overline{\\cal{F}}}_{\\alpha}=0$), and its sign is usually opposite to the sign of $\\alpha_{\\rm K}$ in such a way that the final amplitude of the total $\\alpha$-effect decreases, and so does the final value of the magnetic energy.  \\subsection{Magnetic helicity fluxes}  Recently it has been pointed out that the catastrophic quenching could be alleviated by allowing the flux of small-scale magnetic (or current) helicity out of the domain, so that the total magnetic helicity inside need not be conserved any longer. Alternately, we may introduce those fluxes in the equation for $\\alpha_{\\rm M}$; see equation (\\ref{eq.flux}). Several candidates have been proposed for the helicity fluxes in the past \\citep{kr99,vch01,sb04}. Amongst them are the flux of magnetic helicity across the iso-rotation contours, advective and diffusive fluxes and also the explicit removal of magnetic helicity in processes like coronal mass ejections or galactic fountain flows, for the case of the galactic dynamo.  From the mathematical point of view, the nature of the flux terms in the equation for $\\alpha_{\\rm M}$ has not been demonstrated with sufficient rigor. However, several DNS have pointed to its existence.  Firstly, the shearing box convection simulations of \\cite{kapyla08} showed that in the presence of open boundaries, the large-scale magnetic field grows on temporal scales much shorter than the dissipative time scale. They concluded from this that open boundaries may allow the magnetic helicity to escape out of the system. These experiments seem to be compatible with the flux proposed by \\cite{vch01}, whose functional form may be expressed as \\cite[see][for further details]{sb04,bs05c}: \\begin{equation} \\label{eq.vcf} \\overline{\\cal F}^{\\rm VC}_i = C_{\\rm VC} \\epsilon_{ijl} \\overline{\\mathsf{S}}_{lk} \\overline{B}_j\\overline{B}_k \\quad, \\end{equation} \\noindent where $\\overline{\\mathsf{S}}_{lk}=\\frac{1}{2}(\\overline{U}_{l,k} + \\overline{U}_{k,l})$ is the mean rate of strain tensor and $C_{\\rm VC}$ is a non-dimensional scaling factor. As we assume ${\\bm u}_{\\rm p}=0$, this flux has the following three components: \\begin{eqnarray} \\overline{\\cal F}^{\\rm VC}_r &=& C_{\\rm VC} \\label{eq.vcfr} \\left[\\overline{\\mathsf{S}}_{\\phi r}  B_{\\theta}B_{\\rm r} +  \\overline{\\mathsf{S}}_{\\theta \\phi} (B_{\\theta}^2 - B_{\\phi}^2) \\right]  \\;, \\\\ \\overline{\\cal F}^{\\rm VC}_{\\theta} &=& C_{\\rm VC} \\label{eq.vcft} \\left[- \\overline{\\mathsf{S}}_{\\phi \\theta}  B_{r}B_{\\theta} +  \\overline{\\mathsf{S}}_{r \\phi} (B_{\\phi}^2 - B_r^2) \\right]  \\;, \\\\ \\overline{\\cal F}^{\\rm VC}_{\\phi} &=& C_{\\rm VC} \\label{eq.vcfp} \\left[\\overline{\\mathsf{S}}_{\\theta \\phi}  B_{r}B_{\\phi} -  \\overline{\\mathsf{S}}_{r \\phi} (B_{\\theta}^2 - B_{\\phi}^2) \\right] \\;, \\end{eqnarray} \\noindent with $\\overline{\\mathsf{S}}_{\\phi r} = \\overline{\\mathsf{S}}_{r \\phi} = r \\sin\\theta (\\partial \\Omega / \\partial r)/2$ and $\\overline{\\mathsf{S}}_{\\theta \\phi} = \\overline{\\mathsf{S}}_{\\phi \\theta} = \\sin\\theta (\\partial \\Omega / \\partial \\theta)/2$.  Secondly, \\cite{mitra09}  performed $\\alpha^2$ dynamo simulations driven by forced turbulence in a box with an equator. They found that the diffusive flux of $\\alpha_{\\rm M}$  across the equator can be fitted to a Fickian diffusion law given by, \\begin{equation} \\mbox{\\boldmath ${\\cal F}$}_{\\rm D} = -\\kappa_{\\alpha}(r) \\bm\\nabla \\alpha_{\\rm M} \\; . \\label{eq.fickian} \\end{equation} \\noindent They also computed the numerical value of this diffusion coefficient, and found it to be of the order of turbulent diffusion coefficient. They also found that the time averaged flux is gauge independent. Both results were later corroborated by simulations without equator, but with a decline of kinetic helicity toward the boundaries \\citep{HB10}.  Additionally, magnetic helicity may be advected by the mean velocity with a flux given by $\\bm{\\overline{\\cal F}}_{\\rm ad}=\\alpha_{\\rm M} \\overline{\\bm U} $, or it may be expelled from the solar interior by coronal mass ejections (CMEs) or by the solar wind. This flux, ${\\cal F}_{\\rm CME}$, may account for $\\sim10$\\% of the total helicity generated by the solar differential rotation, as estimated by \\cite{br00}. It can be modeled by artificially removing a small amount of $\\alpha_{\\rm M}$ every $\\tau$ time \\citep{bcc09}, or also by a radial velocity field that mimics the solar wind.  The total flux of magnetic helicity may be written as the sum of these contributions, \\begin{equation} \\bm{\\overline{\\cal F}} = \\bm{\\overline{\\cal F}_{\\rm VC}} + \\bm{\\overline{\\cal F}}_{\\rm D} + \\bm{\\overline{\\cal F}}_{\\rm ad} + \\bm{\\overline{\\cal F}}_{\\rm CME} \\; . \\label{eq.fluxes} \\end{equation} Since in this dynamo model we do not include any component of the velocity field other than the differential rotation, in this study we will consider only the first two terms on the rhs of equation (\\ref{eq.fluxes}). ",
        "conclusions": "\\label{sec.conc} We have developed $\\alpha\\Omega$ dynamo models in spherical geometry with relatively simple profiles of $\\alpha_{\\rm K}$ and shear ($\\partial \\Omega/\\partial r$ and $\\partial \\Omega/\\partial \\theta$). We choose potential field (also vertical field in some cases) and perfect conductor boundary conditions for the top and bottom boundaries, respectively.  We estimate the critical dynamo number by fixing $C_{\\Omega}=-10^4$ and varying $C_{\\alpha}$ while using  algebraic quenching.  Using a dynamo number, $C_{\\Omega}C_{\\alpha}$, that is slightly super-critical, we solve the induction equations for $B$ and $A$ together with an equation for the dynamical evolution of the magnetic $\\alpha$ effect or $\\alpha_{\\rm M}$. We find that for positive (negative) values of $C_{\\alpha}$ in the northern (southern) hemisphere, $\\alpha_{\\rm M}$ is mainly negative (positive), with narrow fractions of opposite sign in regions where $\\alpha_{\\rm K}$ or $\\overline{B}$ are equal to zero.  We find that the kinematic phase is independent of $\\Rm$. However for $\\Rm>10^2$ there exists a phase of relaxation post saturation in which the averaged magnetic field oscillates about a certain mean. The larger the $\\Rm$, the more pronounced are the damped oscillations and the longer is the relaxation time (Fig. \\ref{fig.no.flux}). The final value of the magnetic energy obeys a $\\Rm^{-1}$ dependency ($R_{\\rm m}^{-0.5}$ for magnetic field, Fig. \\ref{fig.brms}), which is in agreement with earlier work \\citep{bs05c,bcc09}.  We argue that including equation (\\ref{eq.am}) in MFD models is appropriate for describing the quenching of the magnetic field in the dynamo process. Since we observe large-scale magnetic fields at high magnetic Reynolds numbers in astrophysical objects, there must exist a mechanism to prevent the magnetic field from catastrophic quenching.  We have studied the role that diffusive and VC fluxes may play in this sense. Their contribution may be summarized as follows: \\begin{enumerate} \\item In the presence of diffusive fluxes, $\\alpha_{\\rm M}$ has only one sign in each hemisphere (negative in the northern hemisphere and positive in southern) and is evenly distributed across the dynamo region (Fig. \\ref{snaps.flux}). \\item For $\\Rm<10^2$ the mean values of $\\alpha_{\\rm M}$ are similar to models without diffusive fluxes, whereas for $\\Rm\\ge10^2$, $\\alpha_{\\rm M}$ has smaller values that seem to be independent of $R_{\\rm m}$ (see Fig. \\ref{fig.brms}, middle). \\item Even a very low diffusion coefficient, e.g.\\ $\\kappa_{\\alpha}=0.001\\etat$, causes $\\overline{B}_{\\rm rms}$ to depart from the $R_{\\rm m}^{-0.5}$ tendency and converge to a constant value which is then around $5$\\% of the equipartition value for large values of $\\Rm$, but below the value of $10^7$ used in this study (dashed line in Fig. \\ref{fig.brms}, top). \\item Larger values of $\\kappa_{\\alpha}$ result in larger final field strengths. \\item In models with only radial shear the Vishniac-Cho flux contributes to $\\alpha_{\\rm M}$ with a component that travels in the same direction as the dynamo wave. This produces a different radial and latitudinal distribution of the magnetic $\\alpha$ effect that also affects the distribution of the magnetic fields. However, it does not help in alleviating the quenching at high $R_{\\rm m}$. On the contrary, the larger the coefficient $C_{\\rm VC}$, the smaller is the resultant magnetic field. \\item In models with only latitudinal shear the VC flux travels radially outward but it remains concentrated at the center of the dynamo region. In a given hemisphere the resultant distribution of $\\alpha_{\\rm M}$ has both positive and negative signs. The part of $\\alpha_{\\rm M}$ that has the same sign as $\\alpha_{\\rm K}$ enhances dynamo action. This effect is more evident in models with vertical field boundary conditions (Figs. \\ref{snaps.vc.flux}b-d). \\item In models with vacuum and vertical field boundary conditions and $R_{\\rm m}=10^3$, the VC flux increases the final value of the magnetic field by a factor of two compared to the case without any fluxes. \\item The magnetic field in models with $R_{\\rm m}\\geq 10^4$ and with non-zero VC flux decays after the kinematic phase since the total $\\alpha$ effect becomes subcritical (see dashed lines in Fig. \\ref{fig.en.vc.flux} a and c). \\item Larger values of $C_{\\rm VC}$ produce narrow bands of $\\alpha_{\\rm M}$ which drives intense dynamo action in these regions. This positive feedback between the magnetic field and $\\alpha_{\\rm M}$ causes the simulation to become numerically unstable in the absence of any other quenching effect. \\end{enumerate}  From the above results it is clear that diffusive fluxes are much more important in alleviating catastrophic quenching when compared to the Vishniac \\& Cho fluxes (in the form of equation \\ref{eq.vcf}) for a large range of $\\Rm$. This is somehow intriguing since it is known from DNS that shear in domains with open boundaries does indeed help in alleviating the catastrophic quenching. It may be understood as a result of the large value of $C_{\\Omega}$ compared with $C_{\\alpha}$ and also to the top boundary condition for the azimuthal magnetic field \\citep{B05,kapyla08}.  The results presented above indicate that considerable work is still necessary in order to understand the role of larger-scale shear in transporting and shedding small-scale magnetic helicity from the domain.  In snapshots of the meridional plane as well as in butterfly diagrams we notice that the diffusive fluxes do not significantly modify the morphology and the distribution of the magnetic field when compared with cases without fluxes or even with simulations with algebraic $\\alpha$ quenching. On the other hand, for models with VC flux the distribution of $\\alpha_{\\rm M}$ becomes different and so does the magnetic field. This is clear from the butterfly diagram shown in Fig.~\\ref{but.vc.flux}b, which exhibits a magnetic field confined to equatorial latitudes reminiscent of the observed butterfly diagram of the solar cycle. Even though this result corresponds to a simplified model, it illustrates the importance of considering the dynamical $\\alpha$ quenching mechanism for modeling the solar dynamo. Similar changes in the distribution of $\\alpha_{\\rm M}$ and $\\bm{\\overline{B}}$ are expected to happen when advection terms are included in the governing equations.  In the simulations presented here, $\\Omega$ and $\\alpha$ effects are present in the same layers. An interesting question is whether the quenching of the dynamo is catastrophic when both layers are segregated, as in the Parker's interface dynamo or the flux-transport dynamo models. We address this question in detail in two companion papers \\citep{cbg10,cgb10}.  We should notice that the back reaction of the magnetic field affects not only the $\\alpha$ effect, but also the other dynamo coefficients, including the turbulent diffusivity. Contrary to quenching of $\\alpha$, the quenching of $\\etat$ may be considered through an algebraic quenching function \\citep[see e.g.][]{ybr03,kb09}. \\cite{gddp09} have shown that in a flux-transport model these effects could affect properties of the models such as the final magnetic field strength and its distribution in radius and latitude. We leave the study of models with simultaneous dynamical $\\alpha$ and $\\eta$ quenchings for a future paper. Solar-like profiles of differential rotation and meridional circulation along with dynamical $\\alpha$ quenching will also be considered in a forthcoming paper."
    },
    "1005/1005.2113_arXiv.txt": {
        "abstract": "{The standard method of mapping the interstellar medium in a galaxy, by observing the molecular gas in the CO 1-0 line and the atomic gas in the 21-cm line, is largely limited with current telescopes to galaxies in the nearby universe. In this letter, we use SPIRE observations of the galaxies M99 and M100 to explore the alternative approach of mapping the interstellar medium using the continuum emission from the dust. We have compared the methods by measuring the relationship between the star-formation rate and the surface density of gas in the galaxies using both methods. We find the two methods give relationships with a similar dispersion, confirming that observing the continuum emission from the dust is a promising method of mapping the interstellar medium in galaxies.}  \\titlerunning{Mapping the ISM with SPIRE}  \\authorrunning{S.A. Eales et al.} ",
        "introduction": "Two of the key measurements of a galaxy that are needed for any investigation of galactic evolution are the star formation rate and, because stars form out of gas, the mass of interstellar material. The standard method of measuring the gas mass is to use the 21-cm line to measure the mass of the atomic gas and the CO 1-0 line to measure the mass of the molecular material. The conversion between 21-cm flux and the mass of atomic gas is unambiguous, but the conversion factor between the CO luminosity and the mass of molecular gas (the notorious `X factor') is uncertain both because the CO 1-0 line is optically thick and because the CO molecule is only tracing the much larger number of unseen hydrogen molecules (e.g. Bell, Viti \\& Williams 2007). Both techniques suffer from the problem that with current telescopes it is very difficult to measure gas masses for galaxies beyond a redshift of $\\simeq$0.2, a major limitation now that the Herschel Space Observatory (Pilbratt et al. 2010) is detecting thousands of galaxies in the high-redshift universe (e.g. Eales et al. 2010).  The idea of using the continuum emission from the dust as an alternative way of mapping the interstellar medium (ISM) has a long history, dating back at least to Hildebrand (1983). More recently, Guelin et al. (1993, 1995) and Boselli, Lequeux \\& Gavazzi (2002) have attempted to determine the X-factor using the continuum emission from the dust and the assumption that the gas-to-dust ratio is the same in both the atomic and molecular phase of the ISM. The problem in attempting to measure the mass of the ISM from the continuum emission of the dust is that the conversion factors between the submillimetre flux and the mass of dust and between the dust mass and the gas mass are both poorly known. These constants have traditionally been estimated from galactic objects, such as reflection nebulae (e.g. Hildebrand 1983), but James et al. (2002) have argued that a better approach is to estimate them from the galaxies themselves, which is the approach we follow here.  In this letter, we have used both methods to map the ISM in two nearby galaxies, M99 and M100, which were two of the first targets to be observed in the Herschel Reference Survey (Boselli et al. 2010), a survey with SPIRE (Griffin et al. 2010) of 323 galaxies in a volume-limited sample of the local universe with distances between 15 and 25 Mpc. To assess the accuracy of the methods, we have used these maps to investigate the relationship between the star-formation rate in these galaxies and the density of the gas out of which the stars form. This is usually paramaterised by a power-law relationship between the star-formation rate per unit volume and the gas density, $\\rho_{SFR} \\propto \\rho_{gas}^N$ (Schmidt 1959), which for a constant scale height will result in a relationship between the star-formation rate per unit area of the disk and the surface density of the ISM which is also a power law with the same index: $\\sum_{SFR} \\propto (\\sum_{gas})^N$. This relationship is of great interest since different theoretical models predict values of $N$ between 0.75 and 2 (Bigiel et al. 2008 and references therein), but in this paper we are less interested in the value of $N$ than in using the fact that there is undoubtedly some intrinsic relationship between $\\sum_{SFR}$ and $\\sum_{gas}$ to assess the errors, both statistical and systematic, in both methods for estimating the surface-density of the ISM. ",
        "conclusions": ""
    },
    "1005/1005.1201.txt": {
        "abstract": "Detection of magnetic-type ($B$-type) polarization in the Cosmic Microwave Background (CMB) radiation plays a crucial role in probing the relic gravitational wave (RGW) background. In this paper, we propose a new method to deconstruct a polarization map on an incomplete sky in real space into purely electric and magnetic polarization type maps, ${\\mathcal{E}}(\\hat{\\gamma})$ and ${\\mathcal{B}}(\\hat{\\gamma})$, respectively. The main properties of our approach are as follows: Firstly, the fields  ${\\mathcal{E}}(\\hat{\\gamma})$ and ${\\mathcal{B}}(\\hat{\\gamma})$ are constructed in real space with a minimal loss of information. This loss of information arises due to the removal of a narrow edge of the constructed map in order to remove various numerical errors, including those arising from finite pixel size. Secondly, this method is fast and can be efficiently applied to high resolution maps due to the use of the fast spherical harmonics transformation. Thirdly, the constructed fields, ${\\mathcal{E}}(\\hat{\\gamma})$ and ${\\mathcal{B}}(\\hat{\\gamma})$, are scalar fields. For this reason various techniques developed to deal with temperature anisotropy maps can be directly applied to analyze these fields. As a concrete example, we construct and analyze an unbiased estimator for the power spectrum of the $B$-mode of polarization $C_{\\ell}^{BB}$. Basing our results on the performance of this estimator, we discuss the RGW detection ability of two future ground-based CMB experiments, QUIET and POLARBEAR. ",
        "introduction": "}  The extreme conditions in the very early Universe produce primordial perturbations of two generic types, namely, density perturbations (scalar perturbations) and relic gravitational waves (tensor perturbations) \\cite{gr74,guth}. In the simplest scenarios, these perturbations are characterized by a nearly scale invariant primordial power spectra. The experimental determination of the parameters specifying these power spectra provides an important method to investigating the physics of the very early Universe. The Cosmic Microwave Background Radiation (CMB) has proved to be a valuable tool in this respect. Scalar and tensor perturbations leave an observable imprint in the temperature and polarization anisotropies of the CMB. The recent experimental effort, including the WMAP satellite \\cite{wmap}, QUaD \\cite{quad}, BICEP \\cite{bicep} and so on \\cite{dasi,cbi,capmap}\\cite{boomerang,maxipol}, has lead to a robust determination of the parameters characterizing the primordial density perturbations. On the other hand, the detection of relic gravitational waves (RGWs) remains an outstanding experimental challenge, and a key task for current, upcoming and future planned CMB observations on the ground \\cite{keating-ground,quad,bicep,clover,POLARBEAR,quiet,quijote,sptpol,actpol}, on balloons \\cite{ebex,pappa,spider} and in the space \\cite{taskforce,planck,b-pol,litebird,cmbpol}.  Both density perturbations and RGWs contribute to the various CMB anisotropy power spectra, namely $TT$ (temperature), $EE$ (``electric\"-type polarization) and $TE$ (temperature-polarization cross correlation) \\cite{cmb1,hucmb,ttee1,ttee2,ttee3,ttee4,ttee5,ttee6,polnarev,ttteee}. In addition, RGWs produce the ``magnetic\"-type ($BB$) polarization that is not produced by density perturbations \\cite{bb}. In principle, all of these information channels ($TT$, $TE$, $EE$ and $BB$) shoul be used to infer the RGW signal in the CMB. However, if the contribution of RGWs to the CMB is small ($r\\lesssim0.05$), the $BB$ channel will the best venue for detecting RGWs \\cite{ttteee}.  The separation of the polarization field into electric and magnetic components is a subtle issue. In practice, in a CMB experiment, one directly observes only the Stokes polarization components $Q$ and $U$. Given a full sky map of the components $Q$ and $U$, one can construct the so-called $E$-mode and $B$-mode of polarization (sometimes referred to as $G$ and $C$ modes, respectively) using spin-spherical harmonics expansion in an unambiguous manner. It is important to keep in mind that, by virtue of construction, $E$ and $B$ modes of polarization are non-local quantities, and require the information on Stokes parameters on the complete sky. However, in realistic cases (ground-based, balloon-borne and satellite experiments), the Stokes parameters are measured only on a fractional portion of the sky. In this situation, the simplest method for constructing electric and magnetic polarization fields, using the spin-spherical harmonics leads to mutual contamination, often referred to as the $EB$-mixing. This $EB$-mixing can become a dominant hinderance for detecting the RGW signal \\cite{challinor}. In order to overcome this difficulty, numerous methods have been developed to separate $E$ and $B$ types of polarization on an incomplete sky \\cite{lewis,pixel,smith,smith2,needlet,cao,kim}. However, these methods suffer from one or several of the following drawbacks - they are slow in practice, they are difficult to realize in pixel space, and/or they lead to partial information loss.    In the current work, we develop a novel method to separate the electric and magnetic components of polarization on a partial sky. In contrast to previous works \\cite{lewis,pixel,smith}, in this paper we focus mainly on the construction of pure $E$ and $B$ types of polarization in the real space, as opposed to constructions in harmonics space. Our method is based on a simple redefinition of electric and magnetic components of polarization, so as to make them only depending on the differential of the $Q$ and $U$ fields. On a full sky, this definition is equivalent to the standard definition. However, due to the differential nature, our definition is directly extendable to polarization maps given on an incomplete sky. The main advantages of our method are as follows. Firstly, the loss of information is small, and is only caused by the removal of a narrow edge around the observed portion of the sky to reduce numerical errors. Secondly, this method is easy to realize for pixelized polarization maps, and is sufficiently fast so as to be practical for high resolution full sky surveys. Thirdly, the method leads to construction of scalar $E$ and $B$ type fields. For this reason, one can directly apply the various techniques developed for temperature anisotropy. In particular, along this route, we construct an unbiased estimator for the $B$-mode of polarization power spectrum and gauge its performance in detecting RGWs.  The outline of the paper is as follows. In Sec.~\\ref{section2}, we introduce the basic notations and explain the general ideas behind our method. In this section we derive the basic equations, and construct the pure electric and magnetic fields ${\\mathcal{E}}(\\hat{\\gamma})$ and ${\\mathcal{B}}(\\hat{\\gamma})$, given the polarization map ($Q(\\hat{\\gamma})$, $U(\\hat{\\gamma})$) on an incomplete sky. In Sec.~\\ref{section3}, we apply this method in practice for a pixelized polarization map. In this section we explicitly construct the ${\\mathcal{B}}(\\hat{\\gamma})$ field, and discuss the various sources of contamination. We go on to discuss the effect edge removal, size of pixels and experimental beam size on the resulting contaminations. In Sec.~\\ref{section4}, we focus on the applying our method to small sky surveys. Combining our method with the pseudo-$C_{\\ell}$ estimator method we construct an unbiased estimator for the $B$-mode power spectrum. We show that, in practice, our estimator performs only slightly worse than an estimator in an idealized situation with no loss of information. Based on this estimator, we analyze the ability to detect relic gravitational waves through their signature in $B$-mode of polarization in two planned ground-based CMB experiments, QUIET and POLARBEAR. We conclude in Sec.~\\ref{section5} with a brief summary of our main results.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  SECTION 2   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "}  In this paper we have proposed a new method to construct pure electric and magnetic type fields ${\\mathcal{E}}(\\hat{\\gamma})$ and ${\\mathcal{B}}(\\hat{\\gamma})$ from polarization field given on an incomplete sky. Due to the differential definitions of these fields, we avoid the so-called $EB$-mixing problem. In practice when working with pixelized maps, residual leakages from various numerical errors require the removal of data from a narrow edge on the boundary of the observed sky. This leads to a minor loss of information in comparison with the idealized lossless case considered in Sec.~\\ref{section4.2a}.  A major advantage of our approach is that the constructed fields ${\\mathcal{E}}(\\hat{\\gamma})$ and ${\\mathcal{B}}(\\hat{\\gamma})$ are scalar. For this reason, the various techniques developed for the analysis of temperature anisotropy maps can be directly applied to these fields. As an important and motivating application, we discuss the construction of an unbiased estimator for the $B$-mode power spectrum $C_{\\ell}^{BB}$, using the pseudo-$C_{\\ell}$ estimator approach. We find that our method is computationally feasible even in the case of high resolution maps. In particular, it takes $2.5$ minutes on a laptop (2.4GHz processor and 2GB memory) to perform all of the calculations, including the the calculation of ${\\mathcal{B}}(\\hat{\\gamma})$ in pixel space with $N_{\\rm side}=512$, and the construction of unbiased estimators for $C_{\\ell}^{BB}$.  With the help of the constructed unbiased estimator, we have investigated the ability to detect gravitational waves through the $B$-mode of polarization in CMB experiment covering $3\\%$ of the sky. In the absence of instrumental noise, we find $S/N=8.26$ for a model with $r=0.1$. This value is $14\\%$ smaller than an idealized situation with no information loss. In the case of realistic experiments, the signal to noise reduces to $S/N=7.05$ for QUIET and $S/N=4.31$ for POLARBEAR.  In conclusion, we would like to point out that, a similar analysis can be applied to large sky surveys. In particular, for Planck satellite and the planned CMBPol experiment, one can construct unbiased estimators for the polarization power spectra $C_{\\ell}^{EE}$ and $C_{\\ell}^{BB}$, by synthesizing the approach outlined in this paper together with the hybrid estimator method suggested in \\cite{hybrid}. We leave this task for future work.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Acknowledgments   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  ~"
    },
    "1005/1005.2325_arXiv.txt": {
        "abstract": "{ We calculate the CMB temperature distortion due to the energetic electrons and positrons produced by dark matter~(DM) annihilation~(Sunyaev-Zel'dovich effect, $\\rm SZ_{\\rm DM}$) in dwarf spheroidal galaxies~(dSphs). In the calculation we have included two important effects which were previously ignored. First we show that the $e^\\pm$ with energy less than $\\sim \\GeV$, which were neglected in previous calculation, could contribute a significant fraction of the total signal. Secondly we also consider the full effects of diffusion loss, which could significantly reduce the density of $e^\\pm$ at the center of cuspy halos. For neutralinos, we find that detecting such kind of SZ effect is beyond the capability of the current or even the next generation experiments, which is consistent with the quantitative description made by S.~Colafrancesco in \\cite{draco_multi}. In the case of light dark matter (LDM) the signal is much larger, but even in this case it is only marginally detectable with the next generation of experiment such as ALMA. We conclude that similar to the case of galaxy clusters, in the dwarf galaxies the $\\rm SZ_{\\rm DM}$ effect is not a strong probe of DM annihilations. } ",
        "introduction": "The physical nature of dark matter (DM) is a great unsolved problem in modern cosmology. A major observational approach to this problem is to look for possible signatures of DM annihilations. In many theoretical models, the DM particles could annihilate and produce $\\gamma$-ray photons and other energetic particles such as electrons and positrons. This is the case, for example, for the supersymmetric DM model \\cite{susyrev}, and the light dark matter (LDM) model \\cite{LDMtheory}. For some early investigations on this subject, see e.g. \\cite{Zeldovich:1980st,Rudaz:1987ry,Ellis:1988qp,Kamionkowski:1990ty,Fargion:1994me}.   Here we consider the energetic electrons and positrons produced by such annihilations. One possible way to reveal or constrain the presence of such energetic electrons and positrons is to look for the Sunyaev-Zel'dovich (SZ) effect induced by them (SZ$_{\\rm DM}$). Energetic electrons or positrons could scatter with cosmic microwave background (CMB) photons, producing a distortion in the CMB spectrum. In clusters of galaxies, the virial temperature is sufficiently high that this effect could be produced by thermal electrons, which is the so called thermal SZ effect \\cite{SZ1}. The bulk movement of electrons with respect to the CMB also cause a similar effect, which is dubbed the kinetic SZ effect \\cite{SZ2}. In galaxies, however, the virial temperature is too low for the thermal electrons to produce significant thermal SZ effect, and the number of free electrons is also too low to produce significant kinetic SZ effect. However, non-thermal energetic electrons and positrons do exist such as cosmic rays in galaxies, which may produce an additional non-thermal SZ effect. In particular, DM annihilation may supply some of such energetic particles. However, DM annihilation is not necessarily the only source of cosmic ray particles. A number of ``normal'' astrophysical sites, such as supernova remnants, AGN jets, pulsar winds etc., may also be partially responsible for cosmic rays. This is why the dwarf spheroidal galaxies ~(dSphs) may be of particular interest to us, as they are DM dominated and also considered to be very inactive, hence observations are less affected by background originating from normal astrophysical processes such as star formation and AGN activity.  The SZ effect induced by a non-thermal distribution of charged particles were studied in \\cite{szn1,szn2}, and then applied to clusters of galaxies~\\cite{sz_cluster1,sz_cluster2,sz_cluster3} and dSphs, where the Draco dwarf was taken as an prime example~\\cite{sz_draco}. These studies suggest that the DM-induced SZ effect could be used to search for signatures of DM annihilation. However, in a companion paper \\cite{Yuan:2009yy} (hereafter Paper I), we calculated the DM-induced SZ effect for clusters. We showed that contrary to previous claims, the DM annihilation induced SZ effect in galaxy clusters is not as large as previously work~\\cite{sz_draco}, and there is little hope to detect it in the foreseeable future, which confirm the quantitative conclusion in \\cite{draco_multi}. Similar conclusion is also derived in a recent study \\cite{2010JCAP...02..005L}. In the present paper, we generalize our earlier work to the case of dSphs, for two DM candidates: the neutralino and the LDM.  The rest of the paper is organized as follows: in Section 2, we discuss our method of calculation. We first briefly review the equations for SZ calculation with a given electron distribution. Then we discuss the density distribution of the DM in dSphs. Lastly the production of electron s and positrons from DM annihilation and propagation of these $e^{\\pm}$ are discussed. This differs somewhat from the case of clusters, as diffusion is more significant. In Section 3, we present the results of our numerical calculation. Finally we discuss our results and conclude in Section 4. ",
        "conclusions": "In this paper, we calculate the non-thermal SZ effect induced by the energetic electrons and positrons produced by the annihilation of neutralino and LDM in dSphs. We take the Draco dwarf as an example to present most of our results, but we have also applied the same calculation to other luminous dSphs in the Local Group.  In our calculation, we obtain the equilibrium $e^{\\pm}$ distribution by fully solving the diffusion equation with the Green's function method. We take synchrotron emission, ICS, ionization loss as well as diffusion loss of $e^{\\pm}$ into account. We find that for small scale systems such as the dwarf galaxies, diffusion effect is crucial. The $e^{\\pm}$ distribution is much less steep than the source function which is proportional to $\\rho_{\\rm DM}^2 $. This important effect was mentioned in \\cite{draco_multi} but was neglected in the numerical calculations in their previous analysis \\cite{sz_draco}, and the resulting $e^{\\pm}$ density could be different by several orders of magnitude. Another difference is that we considered contributions from $e^{\\pm}$ with relatively low energy (i.e. we do not use any arbitrary energy cut off). Such low energy $e^{\\pm}$ could still contribute significantly to the SZ$_{\\rm DM}$ signal.  The SZ effect by the non-thermal $e^{\\pm}$ depends on the density profile of the DM halo. We considered M99, NFW and B95 profiles. As expected, for the cuspy M99 profile the effect is much more significant. Nevertheless, for all profiles we find much smaller signals than previously claimed. Due to the diffusion effect, even for the strongly cuspy M99 profile, the SZ effect predicted for neutralino is too small to be observed with the current or coming generation of instruments. The possible astrophysical foregrounds or contamination may even make the detection more difficult.   Considering that a smaller mass of DM will result in a higher DM number density and annihilation rate, we also investigated the case of LDM, i.e. DM candidate with MeV mass scale. In this case, there is some hope of detecting the DM induced SZ effect, though the effect is also much smaller than previous claims. For the strongly cusped profile assumed, $\\Delta T$ could reach tens of $\\mu$K when the frequency is around or larger than THz. However, we note that THz observation is difficult to do on ground, with the possible exception of a few sites in the Antarctica high plateau.  One of the largest uncertainty in our calculation is the diffusion coefficient of the dSphs, which depends on the magnetic field strength, for which we have no data. We made an assumption on its form and value for the present calculation. It is possible that the magnetic field is much weaker than we assumed, in that case, the SZ effect by non-thermal $e^{\\pm}$ would be even smaller.  These results show that the non-thermal SZ effect induced by DM annihilation in dSphs is small and difficult to observe, even though the expected astrophysical contamination is small. Combined with our earlier results for clusters (Paper I), we conclude that the non-thermal SZ$_{\\rm DM}$ effect is perhaps not a powerful method to detect or constrain DM annihilations."
    },
    "1005/1005.0320_arXiv.txt": {
        "abstract": "{We present deep \\xmm\\ observations and ESO WFI optical imaging of two X-ray-selected fossil group candidates, \\object{\\0216} and \\object{\\2315}. Using the X-ray data, we derive total mass profiles under the hydrostatic equilibrium assumption. The central regions of \\0216\\ are found to be dominated by an X-ray bright AGN, and although we derive a mass profile, uncertainties are large and the constraints are significantly weakened due to the presence of the central source. The total mass profile of \\2315\\ is of high quality, being measured in fifteen bins from $[0.075 - 0.75]\\,R_{500}$ and containing three data points interior to 30 kpc, allowing comprehensive investigation of its properties. We probe several mass models based on the standard NFW profile or on the  S\\'ersic-like model recently suggested by high-resolution N-body simulations. We find that the addition of a stellar component due to the presence of the central galaxy is necessary for a good analytical model fit. In all mass profile models fitted, the mass concentration is not especially high compared to non-fossil systems.  In addition, the modification of the dark matter halo by adiabatic contraction slightly improves the fit. However, our result depends critically on the choice of IMF used to convert galaxy luminosity to mass, which leads to a degeneracy between the central slope of the dark matter profile and the normalisation of the stellar component. While we argue on the basis of the range of $M_*/L_R$ ratios that lower $M_*/L_R$ ratios are preferred on physical grounds and that adiabatic contraction has thus operated in this system, better theoretical and observational convergence on this problem is needed to make further progess.  } ",
        "introduction": "\\label{sec:introduction}  The now well-established cold dark matter (CDM) paradigm lies at the heart of the fiducial scenario for the formation of structure in the Universe. Numerical simulations of hierarchical clustering in the currently-favoured $\\Lambda$CDM cosmology make a number of observationally testable predictions. One example is the existence of a cusped, quasi-universal dark matter density profile that is characterised by the scale radius, $r_s$, at which the logarithmic slope of the profile is -2, and a dimensionless concentration parameter, $c$ \\citep[e.g.,][]{nfw}. The latter parameter exhibits a distinct mass dependence related to the early formation of lower-mass haloes in the hierarchical context \\citep[e.g.,][]{duf08}.  Lying at the nodes of cosmic filaments, clusters and groups of galaxies are dark matter dominated objects whose properties are a sensitive test of these predictions. Indirect evidence for universality in the underlying dark matter distribution in galaxy groups and clusters was indicated by the similarity in X-ray surface brightness and temperature profiles observed with the {\\it ROSAT\\/} and {\\it ASCA} satellites \\citep[e.g.,][]{neu99,vik99,mar98}. However, more detailed, quantitative, information can be gleaned from examination of the total mass profiles of these systems obtained through the hydrostatic equilibrium equation. The vastly improved spatial resolution and throughput afforded by \\xmm\\ and {\\it Chandra} have now allowed detailed investigation of the mass (and thus total density) profiles of moderately large samples of clusters and groups of galaxies \\citep[e.g.,][and references therein]{pa02,pap05,vikh06,hum06,buo07,gasta07,sch07,sun09}. These and other works have confirmed the cusped, quasi-universal nature of the dark matter profile, with a variation in concentration roughly in line with predictions.  However, the presence of a large concentration of baryons in the centre of the system can deepen the potential well and modify the distribution of dark matter. There has been extensive analytical and numerical work on the modification of the mass profile of a dark halo induced by the assembly of baryons in the inner regions. Early results suggested that simple analytical prescriptions based on the conservation of adiabatic invariants gave an accurate description of the halo response. Following the early work by \\citet{bar84}, \\citet{blu86} devised a simple formula to link the dark mass profiles before and after the assembly of a galaxy. Given the initial, spherically-symmetric enclosed mass profiles of the dark matter, $M^{i}_{\\rm{dm}}$(r), and baryons, $M^{i}_b$(r), one may derive the final dark mass profile, $M^f_{\\rm{dm}}$(r), once the final baryonic mass profile, $M^f_b$(r), is specified. The model assumes that dark matter particles move on circular orbits before and after the contraction, and that their initial, $r_i$ , and final, $r_f$ , radii are related by the condition:  \\begin{equation} r_f[M_b^f (r_f)+M_{\\rm{dm}}]=r_i [M_{\\rm{dm}} + M^i_b(r_i)], \\end{equation}  \\noindent where $M_{dm} = M_{\\rm{dm}}(r_f) = M_{\\rm{dm}}(r_i)$ is the dark mass enclosed by each dark matter particle. This modification is generally termed `adiabatic contraction' due to its parameterisation via adiabatic invariants \\citep[e.g.,][]{egg62,blu86,gne04,duf10}. Further work by \\citet{gne04} showed, based on hydrodynamic numerical simulations, that a modification to using the orbit-averaged radius described their results better than the standard circular orbit assumption.  There are few observational tests of this prediction because of the difficulty of obtaining accurate mass profiles this deep into the core of a given object. Stellar dynamical studies suffer from the velocity dispersion anisotropy problem; weak lensing cannot measure the mass on these scales (less than 100 kpc), and strong lensing arcs are preferentially produced in dynamically disturbed systems \\citep[e.g.,][]{bar96} which are unsuitable for the detection of the effect. Although some progress has recently been made on all of these fronts, \\citep[e.g.,][]{san08}, in this context, X-ray observations offer several advantages conducive to investigation of the mass distribution in the central regions. In particular, the overall signal to noise of a good quality X-ray observation can be significant at many hundreds of sigma, and the density squared dependence of the emission results in a centrally-concentrated signal ideal for investigation of the core regions. The only assumption needed is that the intracluster medium (ICM) be in hydrostatic equilibrium.  \\citet{zap06} were the first to test for adiabatic contraction of the dark matter using X-ray observations. Applying several different analytical models (derived from numerical simulations) to the dark matter profile of the $\\sim 3.5$ keV relaxed galaxy cluster \\object{A2589}, they found no evidence for adiabatic contraction. \\citet{gasta07} analysing the mass profiles of 16 galaxy groups, found that addition of adiabatic contraction did not improve the fit. Building on this result, \\citet{hum06} examined the mass profiles of seven early type elliptical galaxies, again finding no strong evidence for adiabatic contraction.  Observational evidence points to the existence of a class of dark matter haloes that have the gravitating mass and hot gas content of groups, but which are dominated in the optical by a single large early-type galaxy \\citep{pon94}. As noted by \\citet{vik99}, such `fossil' systems may have been undisturbed for a very long time if they are the end result of merging within a normal compact group, and thus they may represent the ultimate examples of hydrostatic equilibrium\\footnote{Note that some numerical simulations point to the fossil state being only a transient phase that is terminated by renewed infall from the local environment \\citep{von08}.}. Furthermore, the exceptionally large concentration of baryons in the centre of the potential well makes them ideal systems with which to search for the effect of adiabatic contraction of the dark matter component.  In the present paper we build on our previous work on fossil systems \\citep{zib09b} by investigating the mass profiles of two fossil group candidates, \\0216\\ and \\2315, using X-ray and optical data. The central regions of \\0216\\ are dominated by emission from a bright AGN, and although we recover the total mass profile, the constraints are significantly weakened due to the presence of the AGN. For \\2315, we are able to recover the mass profile to high precision. We find that the central regions are dominated by the stellar mass component, and using the mass profile of the central galaxy from optical data we investigate several parameterisations of the total mass distribution. We find that the best fitting mass parameterisation is obtained when we add a stellar component both in the case of a NFW and a \\citet{nav04} profile, and that the addition of adiabatic contraction slightly improves the fit. However, the result is critically dependent on the assumed IMF used to convert stellar luminosity to mass. We discuss the impact of the assumed IMF on our results, and argue, based on the range of $M_*/L_R$ ratios, that adiabatic contraction has indeed operated in this system.  Unless otherwise noted, we adopt a $\\Lambda$CDM cosmology with $H_0= 70$ km s$^{-1}$ Mpc$^{-1}$ (i.e., $h_{70} = 1$), $\\Omega_M=0.3$ and $\\Omega_\\Lambda=0.7$. All uncertainties are quoted at the 68 per cent confidence level. In the following, $R_\\delta$ is the radius corresponding to a density contrast of  $\\delta$ times the critical density at the redshift of the system. ",
        "conclusions": "We have presented the first deep X-ray and wide field optical imaging observations of two candidate fossil groups, \\0216\\ and \\2315. Based on the criteria established by \\citet{jon03}, and taking into account the various uncertainties involved in the definition of $R_{500}$, we argue that \\0216\\ is a  {\\it bona fide} fossil system and \\2315, if not formally fossil, shares strong physical similarities with this type of object. The X-ray data quality is exceptional for this type of object, extending from $[0.01 - 0.75]\\,R_{500}$ in both cases, allowing us to investigate in detail the properties of their profiles. While unfortunately the central regions of \\0216\\ are contaminated by a bright X-ray point source that contributes $\\sim 40$ per cent of the emission in the central temperature profile bin, we devise a method to correct for its presence and derive the resulting corrected density and temperature profiles.  The object temperatures are $2.05\\pm0.05$ keV and $1.68\\pm0.03$ keV for \\0216\\ and \\2315, respectively, when measured in the $[0.15-1]\\,R_{500}$ region, placing them squarely in the galaxy group category. Both systems exhibit regular, highly peaked X-ray emission centred on the BCG, indicative of their being morphologically relaxed objects. The temperature profiles both describe the shape typical of cool core systems, but with a temperature peak at $\\sim 0.075\\,R_{500}$, closer to the centre than is observed for more massive systems. Their entropy profiles show a considerable excess above the expectations from non-radiative simulations across the entire measured radial range (this is true out to $R_{500}$ for \\0216).  Using the temperature and density profiles, and assuming hydrostatic equilibrium, we calculated the total mass profiles of the two groups. For \\0216, the best fitting NFW model yields $c_{500} = 2.22\\pm0.25$ and $M_{500} = 1.31^{+0.12}_{-0.11} \\times 10^{14}\\ M_\\odot$; however, mass constraints overall are weak due to the uncertainties associated with correction for the AGN, and we do not fit more complex models to these data.  The mass profile of \\2315\\ is of sufficient quality for deeper investigation. We find that consideration of the stellar mass of the central galaxy is essential to provide a good fit to the data. The best fitting mass model is either the S\\'ersic-like profile proposed by \\citet{nav04} with an index $\\alpha$ in agreement with predictions, or an NFW profile, in each case plus a stellar component. The concentration is not especially high compared to non-fossil systems, and appears to be in the range observed and expected for normal systems. Applying adiabatic contraction to the dark matter improves the fit slightly and consistently yields a lower $M_*/L_R$ ratio. Based on the range of derived $M_*/L_R$ ratios and comparison to a range of literature IMFs, we argue that low $M_*/L_R$ fits are preferred on physical grounds, implying that adiabatic contraction has operated in this system. These conclusions are robust to the presence of a possible central X-ray point source.  Clearly the most significant source of uncertainty in our analysis is the IMF. Observation of more extreme fossil systems may allow unambiguous detection of the adiabatic contraction effect on the dark matter, and provide evidence that these systems are indeed older than normal systems. As it stands, in the presence of such excellent data, better observational and theoretical progress on the IMF is necessary to draw definitive conclusions."
    },
    "1005/1005.5169_arXiv.txt": {
        "abstract": "In this paper we present the results of a high resolution (5\\arcsec) CARMA and SZA survey of the 3\\,mm continuum emission from 11 of the brightest (at 1.1\\,mm) starless cores in the Perseus molecular cloud. We detect 2 of the 11 cores, both of which are composed of single structures, and the median 3$\\sigma$ upper limit for the non-detections is 0.2\\solmass in a $\\sim$5\\arcsec\\ beam.  These results are consisent with, and as stringent as, the low detection rate of compact 3\\,mm continuum emission in dense cores in Perseus reported by \\citet{Olmi05}.  From the non-detection of multiple components in any of the eleven cores we conclude that starless core mass functions derived from bolometer maps at resolutions from 10\\arcsec--30\\arcsec\\ (e.g.~with MAMBO, SCUBA or Bolocam) are unlikely to be significantly biased by the blending of lower mass cores with small separations.  These observations provide additional evidence that the majority of starless cores in Perseus have inner density profiles shallower than $r^{-2}$. ",
        "introduction": "Surveys of starless cores in nearby molecular clouds have shown that the starless core mass function (CMF) looks very similar to the stellar initial mass function (IMF), but translated towards higher masses by a factor of a few \\citep[e.g.][]{Motte98, Nutter07, Enoch08, Simpson08}.  If each starless core collapses into a single star, then the apparent shift between the IMF and the CMF implies that the formation process is inefficient and most of the mass in each core never makes it into any star.  Models have shown that 25\\%-70\\% of the mass in a protostellar core will be converted into a star with the remaining material being ejected by outflows \\citep{Matzner00}.  A second reason for the shift between the CMF and the IMF is likely to be fragmentation, in which a single starless core will eventually form multiple stars.  Since many stars are found in binary or higher-order systems \\citep[e.g.][]{Duquennoy91}, to some extent the fragmentation of a starless core into multiple bound objects is to be expected. However, lower mass stars are less frequently found in gravitationally bound systems \\citep{Lada06}.  Although the shape of the IMF can be reproduced from the CMF with a variety of assumptions about the multiplicity, fragmentation, and efficiency of the process \\citep{Swift08, Hachell08}, recent work by \\citet{Goodwin08} has shown that models in which every starless core eventually forms multiple stars reproduces the stellar IMF more successfully than models in which low-mass cores are more likely to form a single star than a binary.  That each core fragments into only two or three components, and not many more, is suggested by the relatively small number of close binaries and high velocity single stars that would result from the dynamical ejection of protostars from higher order systems \\citep{Goodwin05}.  Therefore, high angular resolution observations of starless cores might be expected to show that what appears to be a single object at low resolution is in fact composed of multiple fragments, depending on whether or not the fragmentation takes place before an embedded protostar forms and on what physical scale fragmentation occurs.  The binary fraction of young stellar objects and protostars has been studied in detail, and it has been found that these objects are commonly found in multiple systems.  For example, near-infrared observations of Class-I and flat-spectrum objects in nearby molecular clouds show that in the range 300-2000\\,AU the multiplicity is $\\sim$18\\%, with a median separation between companions of $\\sim$900\\,AU \\citep{Haisch04}.  Sub-arcsecond Berkeley-Illinois-Maryland Array (BIMA) observations of the 3\\,mm continuum emission from 20 Class 0 and Class I protostars found that all of the embedded objects are in small groupings, with separations ranging from $<$100\\,AU to $>$10000\\,AU \\citep{Looney00}.  Recent simulations of fragmentation in magnetically supported dense cores typically start with small cores ($\\sim$6000\\,AU or 24\\arcsec\\ at the distance of Perseus) that fragment into multiple protostellar systems with separations on the order of 10--1000\\,AU \\citep[e.g.][]{Boss02,Boss07,Price07}.  Observations of starless cores have not often found evidence for fragmentation on the scale of 5--10\\arcsec.  In a survey of dense cores in the Perseus molecular cloud with the Owens Valley Radio Observatory (OVRO) interferometer, \\citet{Olmi05} found that five out of six cores had no compact continuum emission, and in the sixth core a single source was detected (at $\\sim$5\\arcsec\\ resolution, or 1200\\,AU in Perseus).  Similarly, \\citet{Harvey03} mapped the 1.3\\,mm continuum emission from the starless core L694-2 with the IRAM Plateau de Bure Interferometer (PdBI) and BIMA and found that the PdBI observations resolved out the core entirely while the shorter baselines of the BIMA data can be adequately modelled by a single source, implying that no fragmentation has taken place.  Other interferometric observations of starless cores, such as S68NW, L1544 and L694-2 show only single peaks in their integrated spectral line emission \\citep{Williams99a, Williams99b, Williams06}.  However, BIMA mapping of the \\nthp\\ emission from the starless core L183 shows that this core is composed of three components with a total mass of a few tenths of a solar mass that are distinct both spatially and kinematically \\citep{Kirk09}.  \\citet{Kirk09} interpret this as evidence that L183 is fragmenting as it collapses and predict that this core will eventually form multiple protostars.  Here we present the results of a search for substructure in a sample of 11 starless cores in the Perseus molecular cloud.  We use Combined Array for Research in Millimeter-wave Astronomy (CARMA) and Sunyaev-Zel'dovich Array (SZA) interferometric observations to study the 3\\,mm continuum emission at $\\sim$5\\arcsec\\ resolution with sensivity up to scales of 30\\arcsec - 80\\arcsec.  The source selection and observations are discussed in Section \\ref{OBSERVATIONS}, followed by an analysis (Section \\ref{ANALYSIS}), a discussion (Section \\ref{DISCUSSION}), and summary (Section \\ref{SUMMARY}). ",
        "conclusions": "\\label{DISCUSSION}  Given that many stars are found in multiple systems \\citep{Duquennoy91} and that many protostars are also multiples \\citep{Looney00, Haisch04}, it is worthwhile to examine when fragmentation takes place.  If starless cores fragment before formation of an embedded object, then we would expect to be able to detect these sub-cores.  On the other hand, if fragmentation takes place after the first protostar has formed, then a survey of starless cores won't find any evidence of fragmentation.  In this study, we find that none of the eleven starless cores show any evidence of fragmentation.  If these cores will eventually fragment, then this must happen during and/or after the collapse and formation of the first protostar.  Late fragmentation on small scales, which is consistent with the results of this survey, is predicted by theories that create binaries from disk fragmentation \\citep[e.g.][]{Bonnell94, Kratter10}.  The observations described in this paper are in agreement with the results of other recent surveys in Perseus. \\citet{Kirk07} compared the line profiles of C$^{18}$O (2-1) and N$_2$H$^+$(1-0) toward a sample of 150 candidate dense cores, both prestellar and protostellar. They found that the N$_2$H$^+$ lines, representing the densest gas within the core, display nearly thermal line widths, indicating little non-thermal activity at the core center. While the C$^{18}$O gas shows significant non-thermal broadening, the velocity centroids of the two species are very similar, indicating that the connection between the cores and their environment is relatively calm. Interestingly, such quiescence across scales appears hard to reconcile with numerical simulations of turbulent clouds \\citep{KirkH09} and may indicate a need for magnetic pressure support within the dense gas. Additionally, \\citet{Jorgensen07} compared the location of mid-infrared {\\it Spitzer} sources and sub-millimeter cores in Perseus, finding excellent agreement between the deeply embedded protostars and their proto-stellar envelopes. Intriguingly, only 3 out of 40 cores were found to harbour multiple protostars, at the spatial resolution of {\\it Spitzer} (MIPS 24\\,\\micron, 6\\arcsec\\ or $\\sim$1500\\,AU). Thus, it appears that within Perseus there is little evidence for fragmentation within cores, at least until the protostellar stage, where the fragmentation length scale is smaller. This may be due to the smooth and quiescent manner in which the cores themselves appear to be forming.  We only detected significant ($\\ge$5$\\sigma$) emission from two of the eleven starless cores in our sample and we put tight upper limits on the masses of multiple components in the other nine cores, from which we conclude that there is no substructure on $\\sim$5\\arcsec\\ scales ($\\sim$1200\\,AU) in the majority of starless cores.  A smooth column density profile is predicted by many models of isolated star formation, such as the inside-out collapse model \\citep{Shu77} and ambipolar diffusion \\citep{Ciolek93}.  It is less clear if a smooth column density profile is consistent with cores created through turbulence, though some simulations show that turbulent molecular clouds do form quiescent cores with smooth column density profiles \\citep{Nakamura05, Ballesteros03}."
    },
    "1005/1005.0785_arXiv.txt": {
        "abstract": "{ We present an X-ray analysis and a model of the nonthermal emission of the pulsar wind nebula (PWN) \\msh. We analyzed \\xmm\\ data to obtain the spatially resolved spectral parameters around the pulsar \\psrb. A steepening of the fitted power-law spectra and decrease in the surface brightness is observed with increasing distance from the pulsar. In the second part of this paper, we introduce a model for the nonthermal emission, based on assuming the ideal magnetohydrodynamic limit. This model is used to constrain the parameters of the termination shock and the bulk velocity of the leptons in the PWN. Our model is able to reproduce the spatial variation of the X-ray spectra. The parameter ranges that we found agree well with the parameter estimates found by other authors with different approaches. In the last part of this paper, we calculate the inverse Compton emission from our model and compare it to the emission detected with the \\hess\\ telescope system. Our model is able to reproduce the flux level observed with \\hess, but not the spectral shape of the observed TeV $\\gamma$-ray emission. } ",
        "introduction": "\\begin{figure} \\centering \\includegraphics[width=\\linewidth]{14151fg01.ps} \\caption{\\xmm\\ count map of \\msh\\ showing the six extraction regions (numbered and marked by white lines) and the exclusion region (black dashed line) used for the spectral analysis. The six extraction regions are centered on the position of the pulsar \\psrb. The exclusion region was chosen to encompass the thermal emission from the \\HII\\ region RCW\\,89.} \\label{msh_skymap} \\end{figure}   \\msh\\ (G320.4$-$1.2) is a complex supernova remnant (SNR) that was discovered in radio wavelengths by \\citet{1961AuJPh..14..497M}. Its radio appearance is dominated by two spots, a brighter one to the northwest and a fainter one to the southeast. The radio emission to the northwest coincides spatially with the \\HII\\ region RCW\\,89. Situated inside the SNR is the energetic $150$\\,ms pulsar \\psrb, which was discovered by the \\einstein\\ satellite \\citep{1982ApJ...256L..45S}. It has a spin-down luminosity of $\\dot{E} = 1.8 \\times 10^{37}\\,\\mathrm{erg}/\\mathrm{s}$ and a characteristic age of $\\tau = 1.7~\\mathrm{kyr}$ \\citep{1994ApJ...422L..83K}. This makes the pulsar one of the youngest and most energetic known. \\citet{1999MNRAS.305..724G} concluded by a comparison of the radio and the X-ray emission that \\msh, \\psrb\\ and RCW\\,89 are associated objects and are located at a distance of $(5.2 \\pm 1.4)\\,\\mathrm{kpc}$.  X-ray observations of the SNR have shown a pulsar wind nebula (PWN) powered by \\psrb\\ \\citep{1996A&A...306..581T,1996PASJ...48L..33T}. Observations with the \\chandra\\ satellite have revealed two outflow jets in the southeast and northwest directions, the latter terminating in the optical nebula RCW\\,89 \\citep[henceforth referred to as G02]{2002ApJ...569..878G}. G02 derived a power-law photon index of $2.05 \\pm 0.04$ and an absorption column density of $(9.5 \\pm 0.3)\\ $\\hcm{21} for the diffuse PWN in the energy band of $0.5$-$10$\\,keV. A detailed study of the innermost region of \\msh\\ was performed by \\citet[henceforth referred to as Y09]{2009PASJ...61..129Y} using an extended data set of \\chandra\\ observations. Their analysis revealed hints for a ring-like feature around the pulsar, which might correspond to a wind termination shock. The hard X-ray spectrum of the PWN was observed with the BeppoSAX and the INTEGRAL satellites. The PWN is clearly seen in the off-pulse component of the emission and the morphology corresponds nicely to the measurements at lower energies. In the energy band of $20$-$200$\\,keV, the off-pulse emission from the PWN is fitted best by a power law with an index of $2.1$ \\citep{2001A&A...380..695M,2006ApJ...651L..45F}.  In the very high-energy (VHE; $100\\,\\mathrm{GeV}$-$100\\,\\mathrm{TeV}$) $\\gamma$-ray domain, the PWN was observed by the High Energy Stereoscopic System (\\hess). The source is clearly extended beyond the point spread function (PSF) and shows a morphology comparable to what is seen in X-rays, extending northwest and southeast of the pulsar. The observed $\\gamma$-ray emission is well-fitted by a power law with a photon index of $2.27$ up to a photon energy of $40\\,\\mathrm{TeV}$\\,\\citep{2005A&A...435L..17A}.  In the first part of this paper, we present an analysis of the \\xmm\\ data of the PWN \\msh. Its large effective area makes \\xmm\\ ideally suited for the spectral analysis of extended sources. Our analysis thus provides a good measurement of the large-scale characteristics of \\msh, compared to earlier high-resolution measurements of the inner region with the \\chandra\\ satellite. Following the \\xmm\\ analysis, we introduce a model which is based on the assumption of the ideal magnetohydrodynamic limit to describe the observed emission in the X-ray domain. Fitting the spatially resolved X-ray emission with the model, we found optimum parameters for several scenarios that have been discussed by other authors. This allows us to constrain physical quantities of the PWN termination shock and the velocity profile of the PWN. In the last part, we also apply this model to make predictions for the VHE $\\gamma$-ray emission as observed by the \\hess\\ Cherenkov telescope array. ",
        "conclusions": "\\label{section_conclusion}  We analyzed \\xmm\\ data of the SNR \\msh\\ to derive the spatially resolved spectral parameters of the inner region of the source. For this analysis we extracted spectra from six annuli centered on \\psrb. A steepening of the X-ray spectrum with increasing distance from the pulsar is observed. The surface brightness in the \\xmm\\ range drops by three orders of magnitude from the inner region close to \\psrb\\ up to the last ring used in our analysis (at a distance of $300$\\,arcsec from the pulsar). The spectra of all rings are fitted well with power laws.  We then introduced a numerical model to describe the spatial evolution of the lepton population and the magnetic field within the PWN. Our model is based on the parameters of the termination shock. We fitted the calculated model spectra to the \\xmm\\ data to constrain the parameters of our model. The results of our optimizations suggest a termination shock in the range of $0.225$ to $0.5$\\,pc (Scenarios~II and III in this work), which are the values derived from different analyses by G02 and Y09. However, the results of the optimizations of our model do not favor one scenario over the other. For the magnetization parameter, we arrive at a lower limit of $0.005$, while our lower limit on the conversion efficiency is $0.3$.  Using the model parameters obtained from the fit to the X-ray data, we calculated the IC emission from the lepton population in the PWN. We rescaled the published \\hess\\ data  to compare the total IC emission from our model with the observed emission. The results show that our model reproduces the right order of magnitude for the observed IC flux, but differs slightly in the spectral shape of the emission. However, we did not take into account contamination effects for the VHE photons due to the PSF. A steeper photon index outside the region used for our model would lead to a systematically steeper spectrum inside the region due to PSF spillover effects. This effect was not taken into account in Fig.~\\ref{fig_SED}. Additionally, we also calculated the spectral index and the surface brightness with increasing distance of the regions from the pulsar, which may be tested with future VHE $\\gamma$-ray instruments like the CTA observatory.  The total energy that is emitted by the leptons in the inner region of the PWN via synchrotron and IC emission in our model can be obtained by integrating over the SED (see Fig.~\\ref{fig_SED}). The integrated emission of $7.6 \\times 10^{35}\\,\\mathrm{erg\\,s}^{-1}$ amounts to roughly $4\\%$ of the current spin-down luminosity of \\psrb. However, when making this comparison one has to be aware that the integrated emission comes from a population of leptons with different lifetimes and cannot be directly linked to the current spin-down luminosity. Furthermore, there is also significant emission at X-ray and TeV energies outside the central \\degree{0.1} region considered in our modeling approach, which increases the percentage of lepton energy converted to radiation (see Figs.~\\ref{msh_skymap} and \\ref{fig_hess_skymap}).  In summary, we show that our radially symmetric model of the PWN \\msh\\ readily describes the nonthermal emission seen in the X-ray band. Despite the simplifying assumptions of our model, the change in spectral index and surface brightness can be reproduced. The model is also able to predict the flux level of the observed emission in VHE $\\gamma$-rays, however, the exact shape of the spectrum is not reproduced. For the future, an application of the model introduced in this paper to other PWN should be interesting to obtain a broader view on the underlying physics in PWN."
    },
    "1005/1005.5396_arXiv.txt": {
        "abstract": "Transit timing variations -- deviations from strict periodicity between successive passages of a transiting planet -- can be used to probe the structure and dynamics of multiple-planet systems. In this paper, we examine prospects for {numerically solving} the so-called inverse problem, the determination of the orbital elements of a perturbing body from the transit timing variations it induces. We assume that the planetary systems under examination have a limited number of Doppler velocity measurements, and show that {a more extensive radial velocity characterization with precision comparable to the semiamplitude of the perturber may remove degeneracies in the solution}. We examine several configurations of interest, including (1) a prototypical non-resonant {system}, modeled after HD40307 b and c, which contains multiple super-Earth mass planets, (2) {a hypothetical system} containing a transiting giant planet with a terrestrial-mass companion trapped in low-order mean motion resonance, and (3) {the HAT-P-13 system}, in which forced precession by an outer perturbing body that is well characterized by Doppler radial velocity measurements can give insight into the interior structure of a perturbing planet, and for which the determination of mutual inclination between the transiting planet and its perturber is a key issue. ",
        "introduction": "{While the overall census of extrasolar planets continues to climb steadily (453 as of this writing\\footnote{{exoplanet.eu, retrieved on May 12, 2010}}),  the emerging population of Earth and Super-Earth sized  ($\\mass \\sin i \\leq 10 \\mass_{\\earth}$) planetary companions that has been uncovered by high-precision radial velocity (RV) surveys \\citep[e.g.][]{Rivera05, Udry07, Mayor09, Vogt10} is shifting the interest of many planet search programs towards terrestrial planets}. Future refinements in ground-based RV programs will likely continue to further push the detection capabilities towards the low-mass end of the planetary population \\citep{Mayor08, Howard10}.  On the other hand, the availability of ground and space-based surveys dedicated to photometric monitoring of large samples of host stars is affording constraints on the true mass and bulk composition of the Super-Earth planetary population \\citep{Leger09, Queloz09, Charbonneau09}.  In particular, the \\Kepler{} mission \\citep[e.g.][]{Koch04, Koch10} is expected to yield transiting Earth-mass planets in the Habitable Zone (HZ) as part of its mission objectives, through continuous and simultaneous photometric sampling of more than 100,000 dwarf stars.  However, this class of objects will likely represent a small percentage of the detections (given the constraints of the mission design), and a large number of Neptune-mass and giant planets will be detected as well \\citep[e.g.][]{Borucki10, Borucki10_4b}.  The exquisite precision and sheer size of the \\Kepler{} transit timing datasets of giant planets, as observed  during the projected four years to six year mission duration, opens up an alternative route to the detection of low-mass planetary companions. Indeed, transit timing variations (TTV) will be caused by gravitational perturbations exerted by additional planets, causing deviations from strictly periodic Keplerian orbits \\citep{MiraldaEscude02, Holman05}. These can be used to infer the orbital elements of the perturbing planet \\citep{Agol05}, or at least place limits on the presence of additional planets \\citep[e.g.][]{Alonso08, MillerRicci08}. An approximate analytic estimate for TTV amplitude for a transiting planet and an external perturber is given by \\citep{Holman05} \\begin{equation} \\delta t \\approx \\ \\frac{45\\pi}{16}\\ \\left(\\frac{\\mass_{pert}}{\\mass_*}\\right) \\ P_{trans}\\ \\alpha_e^3 \\ \\left(1-\\sqrt{2} \\alpha_e^{3/2}\\right)^{-2}\\, \\end{equation}\\label{eqn:ttv} (where we use the symbols $\\mass$ for mass, $P$ for period, $e$ for eccentricity and $a$ for semi-major axis; $\\alpha_e = a_{trans} / \\left[ a_{pert} (1-e_{pert})\\right]$).  The amplitude of these variations can be quite large and amenable to detection, {either in the presence of high-eccentricity perturbers \\citep[e.g.][]{Steffen05}} or when the two planets lie near a low-order mean motion resonance (MMR). Indeed, MMRs are an entirely plausible outcome of core-accretion models of planetary formation, whereby planets can be captured and locked into an MMR during the migration stage \\citep[e.g.][]{Nelson02, Papaloizou05, Beauge06}. Observationally, several of the detected extrasolar systems with multiple planets may be locked in low-order MMRs. {Three such systems (HD82943, HD73526 and HD128311) are engaging in deep 2:1 resonances well characterized by the observations, and GJ876 has recently been reported as a Laplace-type resonance chain \\citep[4:2:1;\\ ][]{Rivera2010}}. {For instance, the TTV amplitude induced by an Earth-mass perturber in a 2:1 resonance with a 3-day Jupiter-mass planet, both in circular orbits, is of order of minutes \\citep{Agol05}}. This is a large signal compared to an accuracy in the measurement of the central transit time of order \\citep{Ford06_2} \\begin{equation} \\sigma_T \\approx \\left(\\frac{t_e}{2\\Gamma}\\right)^{1/2} \\sigma_{ph} \\left(\\frac{R_{pl}}{R_*}\\right)^{-2} \\end{equation} {(where $t_e$ is the duration of the transit ingress/egress, $\\Gamma$ is the observation rate, $\\sigma_{ph}$ is the photometric precision, $R_{pl}$  and $R_*$ are the radius of the planet and the radius of the star, respectively)}, amounting to 10s of seconds for milli-mag photometric accuracy. The recently published \\Kepler{} central times \\citep{Latham10, Borucki10_4b, Jenkins10, Dunham10, Koch10} are in rough accordance with this estimate. Furthermore, with respect to the \\Kepler{} project, we note that once a transit is detected with sufficient signal-to-noise ratio, the star will be switched from the long-cadence (30 minute) to short-cadence (1 minute) sampling rate \\citep{Borucki08}, improving the temporal resolution of the transit even further. We take $\\sigma_{tr, K} = 2\\times 10^{-4}$ d ($\\approx$ 15 s) as a conservative estimate of accuracy on the central transits.  Given a large dataset comprising 1 year or more of continuous transit monitoring, is it possible to infer the mass and elements of the perturbing planet? Reconstructing the properties of the perturber from a noisy TTV signal is a complex, and possibly highly degenerate \\citep{Nesvorny08}, inverse problem. In this paper, we present a series of simulations aimed at detecting low-mass perturbers from realistic central transit and follow-up RV data. To this end, we produce a large sample of \\Kepler-like observations and attempt to characterize the perturber using the algorithm toolset offered by a revised version of the Systemic Console \\citep[][hereafter Paper I]{Meschiari09}.  A number of different planetary realizations were used, in an attempt to fully capture the complexity of TTV fitting, drawing the orbital elements from observed planetary systems. For the sake of simplicity, we focus on two-planet systems, but the method is fully general within the constraints of CPU time and measurement errors.  The plan of the paper is as follows. In \\S \\ref{sec:Console}, we briefly review describe the algorithms used to derive best-fit models and accompanying error estimates. In \\S \\ref{sec:HD40307}, we examine the characterization of planets similar to HD40307\\emph{c} and \\emph{d}, which lie close but not quite in a 2:1 MMR. Our analysis makes use of the HARPS dataset  \\citep{Mayor09} and a simulated transit timing dataset. In \\S \\ref{sec:HATP7} we fit the synthetic realization of a planetary system deep in a 2:1 MMR, with an external perturber (using HAT-P-7 as our model system).  Finally, in \\S \\ref{sec:HATP13} we analyze constraints placed by TTVs on the three-dimensional configuration of planetary systems, using HAT-P-13 as a test case, and conclude in \\S \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  In this paper we outlined a procedure to solve the inverse problem of  deriving best-fitting model parameters and associated uncertainties using synthetic radial velocity and transit timing variations datasets simultaneously. The procedure exploits a number of numerical algorithms that are made available to the community through the Systemic Console package.  We tested our fitting method against a number of synthetic realizations of different planetary configurations, including a system of non-resonating coplanar super-Earths, a system in a deep 2:1  resonance and a non-coplanar system. The transit timing datasets were derived assuming continuous photometric coverage as provided by \\Kepler{}, and thus are fully realistic to the extent that the transit timing error can be modeled as white noise with a constant amplitude. Our analysis shows that combined RV and TTV datasets carry enough dynamical information to characterize a system in its full three-dimensional configuration.  Inverse problems have a storied place in astronomy, with the discovery of Neptune providing a canonical example. In that case, a fortunate orbital geometry allowed Neptune's sky position to be pinpointed with sufficient accuracy that the ``prediction'' of a new planet could credibly be claimed. It is worth pointing out, however, that the accurate ephemeris for Neptune in 1846 was something of a lucky accident. Both Adams' and Le Verrier's masses and semi-major axes were badly off \\citep{Grant1852}. The correct position of the planet that emerged from the calculations stems from a degeneracy of solutions during the period surrounding the conjunction of Uranus and Neptune.  We have found that a similar state of affairs might apply to the transit timing measurement scenarios that will emerge from \\Kepler{}. While departure from strict periodicity can be readily measurable, it is generally difficult to work out the complete system configuration from transit timing measurement alone. We confirmed that the suppression of TTV harmonics by the transit timing noise can lead to severe degeneracies in the model parameters, as first pointed out by \\citet{Nesvorny08}, even when very low levels of timing error is added to the synthetic data.   In presence of such degenerate set of solutions, however, we have verified that adequate RV data can single out the correct orbital configuration. {We note that other constraints derived by extracting more observables from the photometry, such as the the duration of the transits and their variations \\citep[TDV -- ][]{Kipping09, Kipping10}, may also help remove the degeneracies in the solution. Including these contributions will require more sophisticated modelling approaches.}  Finally, we note that our work did not investigate other competing effects that contribute to the TTV signal, chiefly including, but not limited to, light travel time, excitation of tidal modes in the host star, general relativity and the presence of additional planets. Furthermore, the investigation of planetary systems with $N_{pl} >2$ with the methods presented here might be computationally costly due to the large parameter space."
    },
    "1005/1005.5480_arXiv.txt": {
        "abstract": "Data from the \\Fermi-LAT reveal two large gamma-ray bubbles, extending 50 degrees above and below the Galactic center, with a width of about 40 degrees in longitude. The gamma-ray emission associated with these bubbles has a significantly harder spectrum ($dN/dE \\sim E^{-2}$) than the IC emission from electrons in the Galactic disk, or the gamma-rays produced by decay of pions from proton-ISM collisions. There is no significant spatial variation in the spectrum or gamma-ray intensity within the bubbles, or between the north and south bubbles. The bubbles are spatially correlated with the hard-spectrum microwave excess known as the \\WMAP\\ haze; the edges of the bubbles also line up with features in the \\ROSAT\\ X-ray maps at $1.5-2$ keV.  We argue that these Galactic gamma-ray bubbles were most likely created by some large episode of energy injection in the Galactic center, such as past accretion events onto the central massive black hole, or a nuclear starburst in the last $\\sim$ 10 Myr. Dark matter annihilation/decay seems unlikely to generate all the features of the bubbles and the associated signals in \\WMAP\\ and \\ROSAT; the bubbles must be understood in order to use measurements of the diffuse gamma-ray emission in the inner Galaxy as a probe of dark matter physics. Study of the origin and evolution of the bubbles also has the potential to improve our understanding of recent energetic events in the inner Galaxy and the high-latitude cosmic ray population. ",
        "introduction": "\\label{sec:introduction}   The inner Milky Way is home to a massive black hole (MBH), surrounded by clusters of young stars and giant molecular clouds \\citep[see e.g.][for a review]{Morris:1996}.  The nuclear star cluster has a half-light radius of $\\sim 5$pc.  Although there are indications of past activity, the BH is quiescent today.  Fe K$\\alpha$ echoes from molecular clouds around \\sgras\\ have been understood as relics of activity in the past few hundred years \\citep{Sunyaev:1993,Koyama:1996}. On a longer timescale, one might expect relics of past activity in high energy CRs and hot gas, perhaps far off the disk. The most obvious observables would be $e^-$ CR (visible in inverse Compton gammas and microwave synchrotron) and thermal emission (X-rays).  This work presents a multiwavelength study of the inner Galaxy and identifies several large-scale (tens of degrees) gamma-ray features, most notably 2 large (spanning $-50^\\circ < b < 50^\\circ$) structures that we refer to as the ``\\Fermi\\ bubbles''.  We suggest that these bubbles are associated with previously discovered structures in the X-rays and microwaves, and possibly with analogous smaller-scale structures visible in the FIR.  \\subsection{Previous High-Energy Excesses} Observations of gamma-ray emission in the inner Galaxy at $E \\la 1\\gev$ go back decades to \\emph{COS-B} \\citep{Strong:1984cb,COS-B}, and \\emph{SAS-2} \\citep{Fichtel:1975,Kniffen:1981kn} \\citep[see][for a review]{Bloemen:1989}. Later data from the \\emph{EGRET} experiment aboard the \\emph{Compton Gamma-ray Observatory} extended to the high-energy side of the $\\pi^0$ bump\\citep{Smialkowski:1997sm,Dixon:1998di}. However, EGRET lacked the sensitivity and angular resolution to reveal the detailed structure of gamma-ray emission toward the inner Galaxy. The \\emph{Fermi Gamma-ray Space Telescope} provides greatly improved data up to $\\sim100\\gev$, with sufficient angular resolution to map out interesting structures.\\footnote{See \\texttt{http://fermi.gsfc.nasa.gov/ssc/data/}}  At lower energies, the \\emph{ROSAT} All-Sky Survey at 1.5 keV \\citep{Snowden:ROSAT} revealed a biconical X-ray structure over the inner tens of degrees around the Galactic center (GC), later interpreted as a superwind bubble (SWB) with energetics of the order of $10^{54-55}$ ergs \\citep{Sofue:2000, Bland:2003}. On smaller scales, the \\emph{Midcourse Space Experiment} combined with \\emph{IRAS} data also confirms the existence of a limb-brightened bipolar structure, the so called Galactic center lobe (GCL), with origin at the GC on the \\emph{degree} scale \\citep[see e.g.][for a summary of multiwavelength observations of GCL]{Law:2010}. The inferred energy injection of both these bipolar structures, despite their different scales, is $\\sim 10^{54-55}$ ergs, with an estimated age of $\\sim 10^6$ yr for the GCL and $\\sim 10^7$ yr for the SWB. Several Galactic Center Shells (GCS), tens of pc in size, have been found with total energy of order $\\sim10^{51}$ ergs \\citep{Sofue:2003}. These shells and filaments are claimed to originate from one or more episodes of rapid energy release.  \\subsection{Microwave Excess: the WMAP Haze} Beyond direct evidence of shell structures, microwave observations also provide intriguing indications of energy release toward the GC.  At tens of GHz, the {\\it Wilkinson Microwave Anisotropy Probe} (\\WMAP) \\footnote{\\texttt{http://map.gsfc.nasa.gov/}} provides sensitive degree resolution full sky maps of diffuse microwave emission. By subtracting templates including Galactic H$\\alpha$, Haslam 408 MHz soft synchrotron, and thermal dust emission to remove the different known emission mechanisms in these maps, a microwave residual excess (named ``the microwave haze'') with spherical (non-disklike) morphology about $\\sim4$ kpc in radius toward the GC (visible up to at least $|b| \\approx 30\\degree$) has been recognized \\citep{Finkbeiner:2003im}.  It has a spectrum of about $I_\\nu \\sim \\nu^{-0.5}$, harder than typical synchrotron, but softer than free-free. The microwave haze was later interpreted as synchrotron emission from a hard spectrum of $e^-$ cosmic rays. Other hypotheses such as free-free, spinning dust, or thermal dust have failed to explain its morphology, spectrum, or both \\citep{Finkbeiner:2004us,Dobler:2008ww}. However, the most recent \\WMAP\\ 7-year data have not detected the haze polarization predicted by the synchrotron hypothesis \\citep{Gold:2010fm}, implying either heavily tangled magnetic fields, or an alternative emission mechanism.  With that caveat in mind, we will assume the \\WMAP\\ haze is synchrotron and consider the implications.  \\subsection{A Hard Electron CR Spectrum} A simple model, in which the electron CRs that form the haze have diffused from supernova shocks in the disk, cannot fully explain the data for standard diffusion assumptions. The $23-33$ GHz spectrum of the haze synchrotron is as hard as that generated from shocks, and it seems extremely unlikely that these electrons can diffuse several kpc from the disk without significant softening of the spectrum. The synchrotron cooling timescale for cosmic ray electrons emitting at frequency $\\nu$ is $\\tau_{\\rm syn}\\approx10^{6}\\,\\,B_{100}^{-3/2}\\,\\nughz^{-1/2}\\,\\,{\\rm yr}$, where $B_{100} = B/100 \\mu G$ \\citep{Thompson:2006}. Besides the hard spectrum, it is difficult to form the distinctly non-disklike morphology of the haze with any population of sources concentrated in the disk (as is believed to be true of supernovae).  The presence of a distinct component of diffuse hard $e^-$ CR far off the plane is intriguing in itself, and has motivated proposals where the haze is generated by pulsars, other astrophysical processes, or the annihilation of dark matter \\citep{Hooper:2007kb, Cholis:0811, Zhang:2008tb, Harding:2009ye, Kaplinghat:2009ix, McQuinn:2010, Malyshev:2010xc}. Other indications of excess electronic activity in the Milky Way may be found in recent measurements of local electron and positron cosmic rays. The ATIC, \\emph{Fermi} and H.E.S.S experiments have observed a hardening in the $e^+ + e^-$ spectrum at $20-1000$ GeV \\citep{aticlatest, Aharonian:2009ah, Abdo:2009zk}, with an apparent steepening at $\\sim 1$ TeV, and the PAMELA experiment has measured a rising positron fraction above 10 GeV. Taken together, these measurements imply a new source of hard electrons and positrons, which may be related to the \\WMAP\\ haze. The coexistence of \\rosat\\ X-ray bipolar features and the \\WMAP\\ haze toward the inner Galaxy also suggests the interesting possibility of a common physical origin for these signals.  \\subsection{Inverse Compton Excess from Fermi-LAT}  Fortunately, if the \\WMAP\\ haze is synchrotron radiation from a hard electron population located around the GC, the same CRs would also produce IC scattered gammas, allowing an independent probe of the CR population. IC photons provide valuable complementary information about the spatial distribution of the $e^-$ CR (given a model for the ISRF), which in turn can constrain hypotheses about their origin. The unprecedented angular resolution and sensitivity of \\emph{Fermi}-LAT allows us to probe the gamma-ray counterpart to the microwave haze in detail for the first time.  Previous work employing the first year \\Fermi-LAT data isolated a spectrally hard ``gamma-ray haze'' with similar morphology to the \\WMAP\\ microwave haze \\citep{fermihaze}. In this work, we show that the \\Fermi-LAT sky maps constructed from 1.6 yr data (600 days) reveal two large gamma-ray lobes, extending 50 degrees above and below the GC, with a width of about 40 degrees in longitude. These two ``bubble''-like structures have relatively sharp edges and are symmetric with respect to the galactic plane and the minor axis of the galactic disk. The gamma-ray signal reveals similar morphology to the \\WMAP\\ haze, and is also suggestive of a common origin with features in the \\rosat\\ X-ray maps at 1.5 keV towards the GC.  As we will discuss, the sharp edges, bilobular shape, and apparent centering on the GC of these structures suggest that they were created by some large episode of energy injection in the GC, such as a past accretion onto the central black hole, or a nuclear starburst in the last $\\sim$10 Myr. It is well known that the GC hosts a massive black hole and massive clusters of recently formed stars \\citep{Paumard:2006}. Either of these could potentially provide the necessary energy injection by driving large-scale galactic winds or producing energetic jets; we will outline some of the advantages and disadvantages of each scenario.  \\subsection{Structure of This Paper} In \\Refsec{data} we briefly review the \\emph{Fermi}-LAT data and our data analysis procedure. In \\Refsec{bubbles}, we show the 1.6 yr \\Fermi\\ data maps, reveal the \\Fermi\\ bubble features and show that they are robust when different models for the expected Galactic diffuse emission are subtracted. We characterize the morphology of the bubbles in some detail and employ regression template fitting to reveal a hard, spatially uniform spectrum for the gamma-ray emission associated with the bubbles. In \\Refsec{othermaps}, we show that features in the \\rosat\\ soft X-rays and the \\WMAP\\ microwave haze are spatially correlated with the \\Fermi\\ bubbles, and the \\WMAP\\ haze and \\Fermi\\ bubbles are consistent with being produced from the same electron CR population (by synchrotron and IC respectively). \\Refsec{Interp} presents our understanding based on the analysis in \\Refsec{bubbles} and \\Refsec{othermaps}. \\Refsec{explanation} discusses possible scenarios to produce the \\Fermi\\ gamma-ray bubbles. \\Refsec{CR} focuses on the origin of the electron CRs and the challenges in explaining the spectral and spatial profiles of the gamma-ray emission from the bubbles. We discuss the implications of the \\Fermi\\ bubbles for several topics of interest in \\Refsec{discussion}, and give our conclusions and suggest future work in \\Refsec{conclusion}.    \\begin{figure*}[ht] \\begin{center} \\includegraphics[width=1.0\\textwidth]{fermimaps.ps} \\end{center} \\caption{All-sky \\Fermi-LAT 1.6 year maps in 4 energy bins. Point sources have been subtracted, and large sources, including the inner disk ($-2\\degree < b < 2\\degree, -60\\degree < \\ell < 60\\degree$), have been masked. } \\label{fig:fermimaps} \\end{figure*}    \\begin{figure*}[ht] \\begin{center} \\includegraphics[width=1.0\\textwidth]{fermi_model_sub.ps} \\end{center} \\caption{All-sky residual maps after subtracting the \\Fermi\\ diffuse Galactic model from the LAT 1.6 year maps in 4 energy bins (see \\refsec{fermidiffuse}). Two bubble structures extending to $b\\pm 50\\degree$ appear above and below the GC, symmetric about the Galactic plane. } \\label{fig:diffusemodel} \\end{figure*}   \\begin{figure*}[ht] \\begin{center} \\includegraphics[width=1.0\\textwidth]{bubbles_v1.ps} \\end{center} \\caption{ Template decomposition of the \\Fermi-LAT 1.6 year $1-5$ GeV map (see \\refsec{simplediffuse}). \\emph{Top left: }Point source subtracted $1-5$ GeV map, and large sources, including the inner disk $(-2\\degree < b < 2\\degree, -60\\degree < \\ell < 60\\degree)$, have been masked. \\emph{Top middle:} The $1-5$ GeV map minus SFD dust map (\\emph{top right} panel) which is used as a template of $\\pi^0$ gammas. \\emph{Middle row:} The \\emph{left} panel is the same as the \\emph{top middle} panel but stretched $2\\times$ harder. The \\emph{middle} panel subtracts a simple geometric disk template (shown in the \\emph{right} panel), representing mostly inverse Compton emission, to reveal features close to the Galactic center. Two large bubbles are apparent (spanning $-50\\degree < b < 50\\degree$). \\emph{Bottom row:} The \\emph{left} panel is the same as the \\emph{middle} panel of the \\emph{second row}. Finally we subtract a simple bubble template (\\emph{right} panel), with a shape derived from the edges visible in the maps, and uniform projected intensity. After subtracting the bubble template, the two bubbles features have nearly vanished (\\emph{bottom middle} panel), indicating a nearly flat intensity for the \\Fermi\\ bubbles. } \\label{fig:bubbles} \\end{figure*}      \\begin{figure*}[ht] \\begin{center} \\includegraphics[width=1.0\\textwidth]{fullsky_bubble.ps} \\end{center} \\caption{Full sky residual maps after subtracting the SFD dust and disk templates from the \\Fermi-LAT 1.6 year gamma-ray maps in two energy bins. Point sources are subtracted, and large sources, including the inner disk ($-2\\degree < b < 2\\degree, -60\\degree < \\ell < 60\\degree$), have been masked. Two large bubbles are seen (spanning $-50\\degree < b < 50\\degree$) in both cases. \\emph{Right panels:} Apparent \\Fermi\\ bubble features marked in color lines, overplotted on the maps displayed in the \\emph{left panels}. Green dashed circles above and below the Galactic plane indicate the approximate edges of the north and south \\Fermi\\ bubbles respectively. Two blue dashed arcs mark the inner (dimmer) and outer (brighter) edges of the \\emph{northern arc} -- a feature in the northern sky outside the north bubble. The red dotted line approximately marks the edge of \\emph{Loop I}. The purple dot-dashed line indicates a tentatively identified ``donut'' structure. } \\label{fig:fullsky} \\end{figure*}   \\begin{figure*}[ht] \\begin{center} \\includegraphics[width=0.4\\textwidth]{zea_north.ps} \\includegraphics[width=0.4\\textwidth]{zea_north_sub.ps} \\includegraphics[width=0.4\\textwidth]{zea_south.ps} \\includegraphics[width=0.4\\textwidth]{zea_south_sub.ps} \\end{center} \\caption{The zenithal equal area (ZEA) projection with respect to both north pole (\\emph{upper panels}) and south pole (\\emph{lower panels}) for the $1-5$ GeV energy band before (\\emph{left panels}) and after (\\emph{right panels}) subtracting the SFD dust and simple disk templates to reveal the \\Fermi\\ bubbles. } \\label{fig:zea} \\end{figure*}    \\begin{figure*}[ht] \\begin{center} \\includegraphics[width=1.0\\textwidth]{fermiselfclean.ps} \\end{center}  \\caption{\\emph{Top left: } Full sky \\Fermi-LAT 1.6 year $0.5-1.0$ GeV map subtracts the SFD dust map as a template of $\\pi^0$ gammas. \\emph{Top right: } The same as \\emph{top left} panel, but for energy range $2-50$ GeV (note the different gray scale for the two panels). \\emph{Bottom left: }The $2-10$ GeV \\Fermi\\ gamma-ray map subtracting the \\emph{top left} $0.5-1.0$ GeV residual map which is used as a template of ICS of \\emph{starlight}. \\emph{Bottom right:} The same as \\emph{bottom left} panel but for $10-50$ GeV map subtracting the \\emph{top left} $0.5-1.0$ GeV residual map. The \\Fermi\\ bubble structures are better revealed after subtracting the lower energy $0.5-1.0$ GeV residual map with extended disk-like emission. } \\label{fig:fermiselfclean} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusion}  We have identified two large gamma-ray bubbles at $1 \\la E \\la 50\\gev$ in \\Fermi\\ maps containing 1.6 years of data.  They have approximately uniform surface brightness with sharp edges, neither limb brightened nor centrally brightened, and are nearly symmetric about the Galactic plane.  The bubbles extend to $50\\degree$ above and below the Galactic center, with a maximum width of about $40\\degree$ in longitude.  At $|b| \\la 30\\degree$, these ``\\Fermi\\ bubbles'' have a spatial morphology similar to the \\WMAP\\ microwave haze \\citep{Finkbeiner:2003im, Dobler:2008ww}. The \\rosat\\ soft X-ray 1.5 keV map also reveals hard-spectrum features that align well with the edges of the \\Fermi\\ bubbles. The similarities of the morphology and hard spectrum strongly suggest that the \\WMAP\\ haze and the \\Fermi\\ bubbles share a common origin.  In contrast, the \\Fermi\\ bubble features are not aligned with \\emph{Loop I} or any other feature in the Haslam 408 MHz map; while \\emph{Loop I} and other shell structures appear in the gamma rays, their spectra are softer than the bubble spectrum. Furthermore, there are no convincing features spatially correlated with the bubbles in the LAB \\HI\\ or H$\\alpha$ maps.  To better reveal the bubble structures, we use spatial templates to regress out known emission mechanisms. To remove $\\pi^0$ and bremsstrahlung gammas we use the SFD dust map, assuming that the interstellar dust traces gammas produced by cosmic ray protons ($\\pi^0$ decay) and electrons (bremsstrahlung) colliding with the ISM. To trace the gamma rays from IC scattering of disk electrons on the ISRF, we consider 3 different templates: a simple geometric disk model, the 408 MHz Haslam map, and the $0.5-1$ GeV \\Fermi\\ map after removal of the (dust-correlated) $\\pi^0$ gammas.  We verify that our results are insensitive to this choice.  As an additional cross check, we subtract the \\Fermi\\ diffuse Galactic model from the data, finding that this also reveals the bubble structures. Although our template fitting technique is subject to significant uncertainties due to uncertain line of sight gas and CR distributions, these uncertainties mainly affect the intensity profile at low latitudes.  For $|b| > 30\\degree$ and $1 < E < 50\\gev$ the morphology and spectrum are completely consistent for different template choices.  Indeed, in the $1-5\\gev$ maps, a significant part of the southern bubble is easily visible before \\emph{any} template subtractions.  The \\Fermi\\ bubbles have an energy spectrum of $dN/dE \\sim E^{-2}$, significantly harder than other gamma-ray components. There is no apparent spatial variation in the spectrum between the bubble edge and interior, and the north and south bubbles have consistent spatial and spectral profiles. Both the morphology and spectrum are consistent with the two bubbles having the same origin and being the IC counterpart to the electrons which generate the microwave haze seen in \\WMAP\\ \\citep{Finkbeiner:2003im, Dobler:2008ww}. The spectrum of the CR electrons required to generate the \\Fermi\\ bubbles is harder than expected for electrons accelerated in supernova shocks in the disk, and such disk-produced electrons would be even softer after several kpc diffusion; the morphology of the bubble structure is also quite different to that of the lower-energy electrons traced by the Haslam 408 MHz map. Even setting aside the \\WMAP\\ haze, the \\Fermi\\ bubbles are unlikely to originate from excess $\\pi^0$ emission, as (by construction) they are spatially distinct from the SFD dust map, their spectrum is much harder than that of the dust-correlated emission, and the \\rosat\\ data suggest that the bubbles are hot and underdense rather than overdense. The morphology of the \\Fermi\\ bubbles, and the overlap of the \\Fermi\\ bubbles and \\WMAP\\ haze, also strongly \\emph{disfavor} the hypothesis that a significant fraction of the high energy gamma rays observed by \\Fermi\\ in the bubble region are photons directly produced by dark matter annihilation.    The \\Fermi\\ bubble structures were likely created by some large episode of energy injection in the GC, such as a past accretion event onto the central MBH, or a nuclear starburst in the last $\\sim$10 Myr. We have discussed some general possibilities and considerations in this work, and found shortcomings in each scenario; it seems likely that either significant modifications to one of these ideas, or some combination of different mechanisms, will be necessary.  Jets originating from AGN activity can potentially accelerate CR electrons to high energies, and transport them rapidly away from the GC; the cooling time of electrons at $100-1000$ GeV is only $10^{5-6}$ years, so if the CRs are injected and accelerated only in the GC, a very fast bulk transport mechanism is required to convey them throughout the bubbles before they lose a significant fraction of their energy. However, filling the bubbles completely, with $n=10^{-2}$ cm$^{-3}$ gas, would require a mass injection of $\\sim 10^8$ solar masses, so in any case it is more reasonable for the bulk of the material in the bubbles to be swept up and accelerated as the bubbles expand. Energetic shocks associated with jets can have high Mach number and thus efficiently accelerate CR electrons, producing hard spectra with $dN/dE \\sim E^{-2}$, and the total energy required to heat the bubbles is also readily achievable by accretion events onto the central MBH.  However, the north-south symmetry of the bubbles has no obvious explanation in the context of an AGN jet: there is no reason for one jet to be oriented perpendicular to the Galactic plane. The large width and rounded shape of the bubbles are also not typical of jets, which are generally much more collimated, although a precessing jet might help explain the wide opening angle of the bubbles. If the central MBH becomes active on a relatively short timescale, the \\Fermi\\ bubbles may be created by a number of past jets, which combine to give rise to the symmetric and uniform \\Fermi\\ bubbles.  An alternate source for the large required energy injection is a nuclear starburst. The wide opening angle of the bubbles is not a problem in this case; the bubble shape is similar to that observed in NGC 3079, and the X-ray features observed by \\rosat\\ are similar to those observed in other nearby starburst galaxies. However, no corresponding H$\\alpha$ signal of the \\Fermi\\ bubbles is observed, in contrast to other known starburst galaxies: this problem might potentially be resolved if the H$\\alpha$-emitting gas has cooled in the time since the starburst phase (gas hot enough to emit the X-rays observed by \\rosat\\ has a considerably longer cooling time). Also, generally gas filaments and clumps are observed in the X-rays in starburst galaxies, and it would seem that a relic of a past starburst should become more clumpy with gas clouds and filaments due to cooling of the gas. However, while no such structures are obvious in the \\rosat\\ maps, the signal-to-noise is insufficient to place strong constraints.  The absence of any such filamentary structures inside the \\Fermi\\ bubbles, on the other hand, argues against a hadronic origin for the bubble gamma ray emission. Hadronic jets might accelerate protons to high energies, and the interactions of these protons with the ISM could then produce hard $\\pi^0$ gammas and secondary $e^+ e^-$, which would scatter on the ISRF to produce more gamma rays. In this scenario, however, the gamma ray emission should trace the gas density, which we would not expect to be smooth and homogeneous.  Returning to the starburst scenario, the cosmic ray acceleration in this case would be due to shocks at the edge of the bipolar wind. However, the shocks expected in this scenario would be relatively weak and slow-moving, and thus may not be capable of generating a sufficiently hard electron spectrum to reproduce the signal. For example, in first-order Fermi acceleration, a shock Mach number of $\\sim 3.3$ is needed to obtain an electron spectral index of $2.4$, as required for the synchrotron explanation of the \\WMAP\\ haze (see e.g. \\cite{1991SSRv...58..259J}).  The \\Fermi\\ bubbles have sharp edges, also suggesting the presence of a shock at the bubble walls. If the CRs producing the gamma rays have a multi-kpc diffusion length (which is not expected to be the case for 1 TeV electrons, for example), then the edges can still be sharp if the bubble edge is moving outward faster than they can diffuse. If we assume the \\Fermi\\ bubbles are projected structures from three dimensional symmetric blobs towards the GC, the flat intensity profile of the bubbles requires the emissivity to rise at the bubble walls, but remain non-negligible in the bubble interior; the lack of spatial variation in the spectral index may also constrain models where the electrons diffuse long distances from an injection point. Magnetic reconnection in the interior of the bubbles, or some other mechanism such as dark matter annihilation, may help maintain a hard spectrum throughout the bubbles by accelerating existing lower-energy electrons or injecting electrons \\emph{in situ}.  Dark matter annihilation or decay, while an effective mechanism for injecting hard electron CRs at high latitudes, cannot produce the features in the \\rosat\\ X-ray maps correlated with the bubbles, and would not be expected to result in sharp cutoffs in gamma-ray emission at the bubble edges. Dark matter annihilation or decay may be \\emph{contributing} to the bubbles, or to gamma-ray emission in the inner Galaxy that is not well subtracted by either the bubble structure template or the models for known diffuse emission mechanisms; however, understanding the \\Fermi\\ bubbles will be a necessary step before extracting any such dark matter signal.  The \\emph{eROSITA}\\footnote{\\texttt{http://www.mpe.mpg.de/heg/www/Projects/EROSITA/main.html}} and \\emph{Planck}\\footnote{\\texttt{http://www.rssd.esa.int/index.php?project=Planck}} experiments will provide improved measurements of the X-rays and microwaves, respectively, associated with the \\Fermi\\ bubbles, and so may help discriminate between these scenarios. \\emph{eROSITA}, which is expected to launch in 2012, will provide the first imaging all-sky survey of mid-energy X-rays, studying the $0.2-12$ keV energy range with $\\sim 100$ eV energy resolution and a PSF of 20''. The \\emph{Planck} satellite, launched in 2009, will greatly improve the measurements of the \\WMAP\\ haze spectrum.  In addition, AMS-02\\footnote{\\texttt{http://www.ams02.org}} will launch in 2011, and may significantly advance our understanding of CR acceleration and propagation, and help to refine our interpretation of the \\Fermi\\ bubbles.    \\vskip 0.15in {\\bf \\noindent Acknowledgments:} We acknowledge helpful conversations with Avery Broderick, Greg Dobler, Martin Elvis, Jim Gunn, Jill Knapp, Maxim Markivitch, David Merritt, Simona Murgia, Norm Murray, Paul Nulsen, Aneta Siemiginowska, David Spergel, Anatoly Spitkovsky, and Neal Weiner. DF and TS are partially supported by NASA grant NNX10AD85G.  TS is partially supported by a Sir Keith Murdoch Fellowship from the American Australian Association.  This research made use of the NASA Astrophysics Data System (ADS) and the IDL Astronomy User's Library at Goddard.\\footnote{Available at \\texttt{http://idlastro.gsfc.nasa.gov}}  \\newpage"
    },
    "1005/1005.5163_arXiv.txt": {
        "abstract": "Pulsar Timing Arrays (PTAs) use high accuracy timing of a collection of low timing noise pulsars to search for gravitational waves in the microhertz to nanohertz frequency band. The sensitivity of such a PTA depends on (a) the direction of the gravitational wave source, (b) the timing accuracy of the pulsars in the array and (c) how the available observing time is allocated among those pulsars. Here, we present a simple way to calculate the sensitivity of the PTA as a function of direction of a single GW source, based only on the location and root-mean-square residual of the pulsars in the array. We use this calculation to suggest future strategies for the current North American Nanohertz Observatory for Gravitational Waves (NANOGrav) PTA in its goal of detecting single GW sources. We also investigate the affects of an additional pulsar on the array sensitivity, with the goal of suggesting where PTA pulsar searches might be best directed. We demonstrate that, in the case of single GW sources, if we are interested in maximizing the volume of space to which PTAs are sensitive, there exists a slight advantage to finding a new pulsar near where the array is already most sensitive. Further, the study suggests that more observing time should be dedicated to the already low noise pulsars in order to have the greatest positive effect on the PTA sensitivity. We have made a web-based sensitivity mapping tool available at \\url{http://gwastro.psu.edu/ptasm}. ",
        "introduction": "Pulsar Timing Arrays (PTAs) \\citep{1990FosterBacker} are a practical means of observing gravitational waves (GWs) associated with supermassive ($>10^8M_{\\odot}$) black holes, or more generally GWs in the nanohertz to microhertz regime arising from any source \\citep{jb03,jhv+06, 2010SesanaVeccio, Olmez10}. A PTA is a collection of low timing noise pulsars timed to high accuracy. In such a timing array, GWs may signal their presence through correlated disturbances in measured pulse arrival times. Over the last decade, assessments of expected GW source strengths have remained steady \\citep{2010SesanaVeccio} while the number of pulsars that can be timed to high precision has increased, and instrumentation, observing technique, and timing precision of known pulsars have all improved \\citep{Hobbs10}. Detection of GWs via pulsar timing is now within striking distance. Properly made strategic choices for observing and/or improving the timing precision of current array pulsars, and searching for new pulsars, can hasten the arrival of the day when the first detection is made and the field of GW astronomy is properly inaugurated.  Strategic optimization depends upon your strategy. Do you want to detect single sources, or a stochastic background?  Do you want to maximize the number of sources detected, or maximize the signal-to-noise ratio (SNR) of the sources that you do detect? Here we investigate the sensitivity of a PTA to a single source of GWs and investigate optimization strategies that maximize the number of sources detected. We describe the results of a preliminary investigation into strategies for deploying limited observing time among current array pulsars, and for searching for new pulsars to augment the current array.  Our results apply to the search for single sources of GWs\\citep{Finn10, Yardley10}, not to stochastic background searches \\citep{Jenet06, Anholm08, vanHaasteren09}.  After an introduction of the PTA single-source sensitivity calculation and NANOGrav characteristics in \\S 2, we discuss in \\S 3 the consequences of adding pulsars to the array. Further, discussed in \\S 4, is the preliminary study of pulsar timing allocation time. Finally, in \\S 5, we discuss the implications that rise from this study and suggest future research. ",
        "conclusions": "Given the goal to directly detect GWs, the optimization of pulsar timing arrays should be well understood. Figure~\\ref{fig:AddWhereComp} suggests that for the sake of GW detection we may not need to ``fill in\" regions of the sky currently devoid of PTA pulsars, but that rather a clustering of good pulsars in one region yields the gratest number of detectable GW sources. The figure shows the sensitivity (volumetric) gain produced by adding a new pulsar to the array as a function of the location of the added pulsar. The current pulsars in the array are clustered around 18h right ascension and 0 declination (as shown by green diamonds on the figure) and in fact the greatest improvement in the volume to which the array is sensitive is produced by adding a new pulsar near that already existing cluster of pulsars, although the range of improvement is modest in all cases (from improving it not at all to improving it by 6\\%). Here we suggest that if the goal is to maximize the volume of space to which we are sensitive to GW sources the location with highest concentration of pulsars is slightly favored over other locations, but does not make a significant difference.  So perhaps efforts to ``fill in the gaps'' in the spatial arrangements of PTAs are unfounded. However, any new pulsar, regardless of its location, is beneficial to the scientific community. Continuing research is needed on this particular area to optimize pulsar searches for PTA goals.  Our initial findings indicate that the sensitivity to burst and continuous GWs can be significantly improved by observing longer the pulsars for which we already have good timing values. In other words, it is found that the intuitive thing to do, observing longer the pulsars that we do not have good timing values for, significantly decreases our sensitivity to GWs by almost 50\\% (Figure ~\\ref{fig:ImproveWhat}). This suggestion taken to an extreme yields the ridiculous result that it is best to spend all observing time on a signal pulsar.  The sensitivity plot in this case would be the familiar beam pattern of a single pulsar, but if somehow GWs could be detected with a single pulsar, spending all our time on this one best pulsar, assuming the RMS reduces as the square-root of observing time, would in fact maximize the volume of space to which we are sensitive. We discussed in \\S \\ref{add_a_pulsar} that this strategy neglects issues with confirming the detection as distinct from clock, ephemerides, and other errors, so more work must be done to optimize observing strategies, but this work clearly indicates that convention may not be the best means to a discovery. In particular, as the PTA is being optimized, what should the observing strategy be? Is the answer different depending on whether our aim is to detect single sources or a stochastic background. Presented here is a crucial first step toward answering these questions."
    },
    "1005/1005.4336.txt": {
        "abstract": "{} % aims heading (mandatory) {Our purpose is to provide reliable stellar parameters for a significant sample of eclipsing binaries, which are representative of a whole dwarf and metal-poor galaxy. We also aim at providing a new estimate of the mean distance to the SMC and of its depth along the line of sight for the observed field of view.} % methods heading (mandatory) {We use radial velocity curves obtained with the ESO FLAMES facility at the VLT and light curves from the OGLE-II photometric survey. The radial velocities were obtained by least-squares fits of the observed spectra to synthetic ones, excluding the hydrogen Balmer lines.} % results heading (mandatory) {Our sample contains 23 detached, 9 semi-detached and 1 overcontact systems. Most detached systems have properties consistent with stellar evolution calculations from single-star models at the standard SMC metallicity $Z = 0.004$, though they tend to be slightly overluminous. The few exceptions are probably due to third light contribution or insufficient signal-to-noise ratio. The mass ratios are consistent with a flat distribution, both for detached and semi-detached/contact binaries. A mass-luminosity relation valid from $\\sim$4 to $\\sim$18 $\\mathcal{M}_{\\odot}$ is derived. The uncertainties are in the $\\pm$2 to $\\pm$11 \\% range for the masses, in the $\\pm$2 to $\\pm$5 \\% range for the radii and in the $\\pm$1 to $\\pm$6 \\% range for the effective temperatures. The average distance modulus is $19.11\\pm0.03$ ($66.4\\pm 0.9$ kpc). The moduli derived from the $V$ and from the $I$ data are consistent within $0.01$~mag. The $2\\,\\sigma$ depth of the SMC is, for our field, of $0.25$~mag or $7.6$~kpc under the assumption of a gaussian distribution of stars along the line of sight. Three systems show significant apsidal motion, one of them with an apsidal period of 7.6 years, the shortest known to date for a detached system with main sequence stars.} % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "Since the late 1990s, the usefulness of extragalactic eclipsing binaries has been emphasized in a number of papers. The reader can notably refer to the excellent reviews from Clausen (\\cite{jC04}) and Guinan (\\cite{eG04}, \\cite{eG07}). The two major contributions of eclipsing binaries (hereafter EBs) to astrophysics are to provide (1) fundamental mass and radius measurements for the component stars, allowing to test stellar evolution models, and (2) precise distance moduli (DMi) derived from the luminosities calculated from the combination of the absolute radii with the effective temperatures. Until a purely geometrical distance determination is feasible, Paczy\\'{n}ski (\\cite{bP01}) considers that detached EBs are the most promising distance indicators to the Magellanic Clouds. Besides, Wyithe \\& Wilson (\\cite{WW02}, hereafter WW02) remarked that semi-detached EBs are even more promising, since their parameters are better constrained.  The renewal of interest in extragalactic EBs, especially EBs in the Magellanic Clouds, has been stimulated by the release of a huge number of light curves as a byproduct of automated microlensing surveys (EROS, MACHO, OGLE) with 1-m class telescopes. As photometry is only half of the story, high resolution spectrographs attached to 4-m class or larger telescopes had to be used to obtain reliable radial velocity (RV) curves. Four B-type EB systems belonging to the Large Magellanic Cloud (LMC) were accurately characterized in a series of papers from Guinan et al. (\\cite{GFD98}), Ribas et al. (\\cite{RGF00}, \\cite{RFMGU02}) and Fitzpatrick et al. (\\cite{FRG02}, \\cite{FRGMC03}). More recently, from high resolution, high $S/N$ spectra obtained with UVES at the ESO VLT, the analysis of eight more LMC systems was presented by Gonz\\'{a}lez et al. (\\cite{GOMM05}). A corner stone in the study of EBs in the SMC was set up with the release of two papers from Harries et al. (\\cite{HHH03}, hereinafter HHH03) and Hilditch et al. (\\cite{HHH05}, hereinafter HHH05) giving the fundamental parameters of a total of 50 EB systems of spectral types O and B. The spectroscopic data were obtained with the 2dF multi-object spectrograph on the 3.9-m Anglo-Australian Telescope. To our knowledge this is the first use of multi-object spectroscopy in the field of extragalactic EBs. Recently, even the distance to large spiral galaxies were measured on the basis of two EBs in M31 (Ribas et al. \\cite{RJVFHG05}, Vilardell et al. \\cite{Vil10}) and one in M33 (Bonanos et al. \\cite{BSK06}).  The huge asymmetry, between the number of EB light curves published so far and the very small number of RV curves, is striking. If one considers the SMC, the new OGLE-II catalog of EBs in the SMC (Wyrzykowski et al. 2004) contains 1350 light curves and currently only 50 of these systems have moderately reliable RV curves. This paper reduces a little this imbalance by releasing the analysis of 28 more EB systems plus revised solutions for 5 systems previously described by HHH03 and HHH05. The RV measurements were derived from muti-object spectroscopic observations made with the VLT FLAMES facility.  Another strong motivation for increasing the number of fully resolved binaries is to settle the problem of the distribution of the mass ratio of detached binaries with early B primaries. Recently, two papers were",
        "conclusions": "The first goal of the present paper was to extract the fundamental stellar parameters characterizing 33 SB2 systems and then to test the reliability of stellar evolution models at low metallicity. We have shown that most components of detached systems have properties in fair agreement with the predictions of the Geneva single-star models at $Z =0.004$, although several systems appear slightly overluminous for their mass. From the subsample of detached systems, we give a mass-luminosity relation valid for masses between $\\sim 4$ and $\\sim 18 \\mathcal{M}_{\\odot}$. Despite the lack of evolutionary cross checks, the semi-detached binaries appear to satisfy the typical morphology of Algol-like systems: a main-sequence primary with a large, less massive and more evolved secondary companion.  Our binaries were then used as primary standard candles. Average distance moduli obtained from the $I-$band data show no systematic difference compared to values obtained from the $V$ data. For each band, the frequency distribution of the DMi is not significantly different from a Gaussian. Combining $V$ and $I$ results and considering both the mean and the median values, we adopt $\\langle(m-M)_0\\rangle =19.11\\pm 0.03$ as the average DM for the 19 most reliable systems. We have shown that, even if the effective temperatures are overestimated, so as to make the components of our systems overluminous relative to their mass, the average DM would still remain above $19.00$, so that our results support rather the long distance scale to the Magellanic Clouds than the short one. The true depth of our sample is estimated to be $2\\times 0.126=0.25~\\mathrm{mag}$ ($\\sim$7.6 kpc), if we consider the $2\\;\\sigma$ interval as representative.  We have found a significant apsidal motion for at least three systems. For one system, the apsidal motion might be affected by an unseen third companion, but at least two other systems most probably owe it to the stellar structure of their components alone. The observed and predicted precession rates agree perfectly for the latter two systems, which are now circularizing their orbit, and the agreement extends to one or two other systems showing much less precise apsidal motion. The system 5 266131 shows probably the largest precession motion of all systems hosting main sequence stars, which is due exclusively to the non-point mass nature of the components.  Beside the astrophysical parameters and the distance moduli, other interesting results have been obtained. We have shown that the twin hypothesis for massive binaries is not supported by our results, although the latter are not strong enough to reject it either. Our sample of 23 detached systems agree quite well with a standard flat mass ratio spectrum. There is a clear cut-off of the $q-$distribution at $\\sim0.72$. Therefore, the high proportion of detached binaries having a $q$ value close to unity is more probably the result of an observational bias. The discrepancy observed by Pinnsonneault \\& Stanek (\\cite{PS06}) may well be due to a number of underestimated mass ratios in the HHH03/05 sample. The distribution of the mass ratios of the semi-detached/contact binaries is consistent with a flat distribution restricted to the interval $0.4 < q < 0.8$. The sample is too small (10 objects) to tell more.  Another important result is the bimodal distribution of the radial velocity of the \\ion{H}{ii} gas, with peak values at $145.2\\pm 0.5$ and $161.4\\pm 0.6$ km s$^{-1}$.  Finally, we must emphasize that the accomplishment of this project was possible thanks to the public availability of the OGLE-II catalog of EB light curves, the OSTAR2002/BSTAR2006 libraries of synthetic spectra, the KOREL code and the PHOEBE interface to the WD codes."
    },
    "1005/1005.3329_arXiv.txt": {
        "abstract": "We present a new, BVI$_{c}$ photometric survey of the young open cluster IC~4665, which improves on previous studies of this young cluster by incorporating a rigorous standardization procedure, thus providing high-fidelity colors and magnitudes for cluster members. We use this new photometric dataset to reevaluate the properties (age and distance) of IC~4665. Namely, using a statistical approach incorporating $\\tau^{2}$ CMD modeling, we measure a pre-main-sequence isochrone age and distance of 36$\\pm$9 Myr and 360$\\pm$12 pc, as well as a upper-main-sequence {\\em turn-off} age and distance of 42$\\pm$12 Myr and 357$\\pm$12 pc. These ages and distances are highly dependent on the isochrone model and color used for the fitting procedure, with a possible range of $\\sim$10-20 Myr in age and $\\sim$20 parsecs in distance. This spread in calculated ages and distances seen between colors and models is likely due to limitations in the individual membership catalogs and/or systematic differences in the predicted stellar parameters from the different sets of models. Interestingly, when we compare the isochrone ages for IC~4665 to the published lithium depletion boundary age, 28$\\pm$5 Myr, we observe that this cluster does not appear to follow the trend of isochrone ages being 1.5 times smaller than lithium depletion boundary ages. In addition, comparing the overall magnetic activity (X-ray and H$\\alpha$ emission) in IC~4665 with other well studied open clusters, we find the observed activity distributions for this young cluster are best characterized by assuming an age of 30--40 Myr, thus in agreement with our pre-main-sequence and {\\em turn-off} isochrone ages for IC~4665. Overall, although some age discrepancies do exist, particularly in the ages measured from pre-main-sequence isochrones, the range of possible IC~4665 ages derived from the various dating techniques employed here is relatively small compared to that found for other well studied open clusters. ",
        "introduction": "Open clusters have long provided necessary empirical constraints for our current models of stellar evolution due to their large collection of coeval, equidistant, and chemically homogeneous stars. Through observations of large numbers of open clusters with known ages, we are able to investigate the manner with which fundamental stellar properties evolve as stellar evolution progresses. This is particularly true for early stellar evolution (stellar ages $<$ 100 Myr), when fundamental physical parameters of stars change rapidly ({\\em e.g.,} radius, temperature, rotation rate, magnetic field strength). However, for these clusters to be used as valuable fiducial benchmarks of stellar evolution, we must be able to date them in a consistent manner with high-fidelity and precision. Currently, there are several dating methods utilized to determine open cluster ages, including fitting theoretical model isochrones to empirical H-R diagrams, light element abundance distributions, and exploiting the connection between magnetic activity and age.  A heavily relied upon approach to determining absolute ages of star clusters is finding a best-fitting color-magnitude diagram (CMD) isochrone for the nuclear turn-off \\citep{Sandage1958,Mermilliod1981a,Mermilliod1981b,Meynet1993} and/or the pre-main-sequence (PMS) turn-on \\citep{DAntona1997,Baraffe1998,Siess2000} to the cluster's main-sequence, typically performed using an {\\em ad-hoc} `by-eye' comparison of the model isochrone to the observational photometric data. This technique, however easily applied in theory, is in practice commonly hampered by uncertainties including field star contamination, cluster distance and/or photometric error. Furthermore, this technique suffers from the measured cluster ages being heavily dependent on the input physics used by the stellar evolution models. For instance, one can show for the same cluster that different nuclear turn-off and PMS models can give stellar ages that differ by a factor of 1.5-2 \\citep{Chiosi1992,Meynet1993,Meynet1997,Meynet2000,Baraffe1998,Baraffe2002,Siess2000,Naylor2009}.  There now exist several alternative dating techniques that employ other observable properties of stars to find absolute and relative ages for open clusters. Observations have shown that the level of magnetic activity, indicated by the presence of X-ray and/or chromospheric activity, in solar-type stars clearly decreases with increasing age \\citep{Wilson1963,Skumanich1972,Noyes1984,Soderblom1991,Mamajek2008,Cargile2009}. This relationship is fundamentally tied to the interplay between magnetic activity and the rotationally-induced, magneto-hydrodynamic dynamo in solar-type stars. This causal relationship, the so called age-rotation-activity paradigm, generally holds true; however, it can be affected by natural activity cycles, similar to those seen in the Sun. For instance, \\citet{Soderblom1991} show that uncertainties in stellar chromospheric ages can be as large as $50\\%$ for individual stars. Of course, in open clusters this error can be statistically reduced by measuring many stars.  IC~4665 is currently one of only a handful of well-studied PMS open clusters ({\\em i.e.,} age of 5-50 Myr), {\\em e.g.,} NGC~2169 \\citep[10 Myr][]{Jeffries2007b}, NGC~2547 \\citep[35 Myr][]{Jeffries2005}, and IC~2391 \\citep[50 Myr][]{BarradoyNavascues2004}. IC~4665 thus offers an interesting, additional opportunity to place empirical constraints on models of stellar evolution at young ages. Because of this, IC~4665 has been surveyed extensively to determine cluster membership, as well as observe physical properties of cluster members, including studies on the kinematics of the cluster \\citep{Prosser1993,Prosser1994,Manzi2008,Jeffries2009}, a ROSAT survey of stellar X-ray emission \\citep{Giampapa1998}, and studies of the distribution of Li abundances in the cluster \\citep{Martin1997,Manzi2008,Jeffries2009}. Most of these earlier investigations into the properties of IC~4665 members relied on the optical photometry and membership catalogs reported in \\citet[P93]{Prosser1993}. Unfortunately, the cluster members detailed in this study rely upon uncertain photometric data, due to the use of an improper photometric standardization procedure (see Section~\\ref{sec.comp}), as well as having a high level of field star contamination as a result of their membership catalog being based upon unreliable proper motions (see Section \\ref{sec.lowmass}).  In light of such uncertainties, we report a new, standardized BVI$_{c}$ photometric survey of the central 1 sq.~deg.~of the cluster (see Section \\ref{sec.newphot}). We use this new catalog to re-derive the age of IC~4665 using several different methods, including modeling the PMS (Section \\ref{sec.lowmass}) and upper-main-sequence (Section \\ref{sec.highmass}) photometric membership, and comparing the X-ray emission (Section \\ref{sec.xray}) and H$\\alpha$ emission (Section \\ref{sec.halpha}) distributions to other well known open clusters. In addition, we compare our new age measurements in the context of the available ages for IC~4665 in the literature. ",
        "conclusions": ""
    },
    "1005/1005.3059_arXiv.txt": {
        "abstract": "{The origin of the far-infrared emission from the nearby radio galaxy M87 remains a matter of debate. Some studies find evidence of a far-infrared excess due to thermal dust emission, whereas others propose that the far-infrared emission can be explained by synchrotron emission without the need for an additional dust emission component. We present {\\em{Herschel}} PACS and SPIRE observations of M87, taken as part of the science demonstration phase observations of the Herschel Virgo Cluster Survey. We compare these data with a synchrotron model based on mid-infrared, far-infrared, submm and radio data from the literature to investigate the origin of the far-infrared emission. Both the integrated SED and the {\\em{Herschel}} surface brightness maps are adequately explained by synchrotron emission. At odds with previous claims, we find no evidence of a diffuse dust component in M87, which is not unexpected in the harsh X-ray environment of this radio galaxy sitting at the core of the Virgo Cluster.} ",
        "introduction": "At a distance of 16.7~Mpc \\citep{2007ApJ...655..144M}, M87 is the dominant galaxy of the Virgo Cluster. It is one of the nearest radio galaxies and was the first extragalactic X-ray source to be identified. Because of its proximity, many interesting astrophysical phenomena can be studied in more detail in M87 than in other comparable objects (see e.g.\\ \\citet{1999LNP...530.....R} for an overview). Among its many remarkable features is the several billion solar mass supermassive black hole at its centre \\citep{1997ApJ...489..579M, 2009ApJ...700.1690G} and the prominent jet extending from the nucleus, visible throughout the electromagnetic spectrum. The central regions of M87, in particular the structure of the jet, have been studied and compared intensively at radio, optical, and X-ray wavelengths \\citep[e.g.][]{1991AJ....101.1632B, 1996A+A...307...61M, 2001A+A...365L.181B, 2001ApJ...551..206P, 2004ApJ...607..294S, 2005ApJ...627..140P, 2007ApJ...668L..27K, 2008A+A...482...97S, 2010arXiv1003.5334W}.  Compared to the available information at these wavelengths, our knowledge of M87 at far-infrared (FIR) wavelengths is rather poor. A controversial issue is the origin of the FIR emission in M87, i.e., the question of whether the FIR emission is caused entirely by synchrotron emission or whether there is an additional contribution from dust associated with either the global interstellar medium or a nuclear dust component. This question is partly driven by the observation of faint dust features in deep optical images \\citep{1993ApJ...413..531S, 2006ApJS..164..334F}. Several papers on the FIR emission of M87 arrive at different conclusions. \\citet{2007ApJ...663..808P} present ground-based Subaru and {\\em{Spitzer}} IRS spectra of the M87 nucleus and find evidence of an excess at wavelengths longer than 25~$\\mu$m, which they attribute to thermal emission from cool dust at a characteristic temperature of some 55~K. This claim is countered by \\citet{2009ApJ...705..356B}, who present a higher signal-to-noise IRS spectrum of the nucleus. After careful subtraction of a stellar emission template from the mid-infrared spectrum, these authors conclude that the nuclear spectrum can be fully accounted for by optically thin synchrotron emission and that there is little room for dust emission. On a larger scale, \\citet{2004A+A...416...41X} present ISOCAM imaging of M87 and argue that the mid-infrared flux can be attributed to a single synchrotron emission component. \\citet{2007ApJ...655..781S} present {\\em{Spitzer}} IRAC and MIPS imaging of M87 and find a slight excess in the FIR over a power-law interpolation. They attribute this excess emission to dust emission from the host galaxy. Finally, \\citet{2008ApJ...689..775T} observed 1.3~mm continuum emission from the nucleus and jet of M87 and found that the measured fluxes are generally consistent with synchrotron emission, although they could not rule out a possible nuclear contribution from thermal dust emission.  The recently launched {\\em{Herschel}} Space Observatory \\citep{Herschel} offers the possibility to study M87 at FIR wavelengths in more detail than has been possible to date. The PACS \\citep{PACS} and SPIRE \\citep{SPIRE} instruments combined can produce images over the wavelength range between 70 and 500~$\\mu$m with unprecedented sensitivity and superior resolution. In this Letter, we present PACS and SPIRE imaging of M87 at 100, 160, 250, 350, and 500~$\\mu$m, taken as part of the science demonstration phase (SDP) observations of the Herschel Virgo Cluster Survey (HeViCS\\footnote{More details on HeViCS can be found on \\href{http://www.hevics.org}{http://www.hevics.org}.}). We combine these observations with mid-infrared, FIR, submm, and radio data from the literature to investigate the level and the origin of the FIR emission in M87. In Sect.~{\\ref{Observations.sec}}, we present the observations and data reduction, Sect.~{\\ref{Analysis.sec}} describes the analysis of the data and in Sect.~{\\ref{Conclusion.sec}} we present our conclusions. ",
        "conclusions": "\\label{Conclusion.sec}  For both the integrated SED and the SPIRE 250~$\\mu$m map, we have found that synchrotron emission is an adequate explanation of the FIR emission. We do not detect a FIR excess that cannot be explained by the synchrotron model. In particular, we have no reason to invoke the presence of smooth dust emission associated with the galaxy interstellar medium, as advocated by \\citet{2007ApJ...655..781S}. For a dust temperature of 23~K, which is the expected equilibrium temperature in the interstellar radiation field of M87, we find an upper limit to the dust mass of $7\\times10^4~M_\\odot$. Our result agrees with the analysis of the nuclear emission by \\citet{2009ApJ...705..356B}. \\citet{Clemens} discuss the lifetimes of interstellar dust grains in elliptical galaxies in the Virgo Cluster based on {\\em{Herschel}} observations and find an upper limit to the amorphous silicate grain survival time of less than 46 million years. Given that M87 is a luminous X-ray source, the absence of a substantial dust component is not a surprise. A low dust content is also in agreement with the non-detection of cool molecular gas \\citep{2008A+A...489..101S, 2008ApJ...689..775T} and the non-detection of significant intrinsic absorption in the X-ray spectra of M87 \\citep{2002A+A...382..804B}. Our conclusion is that, seen from the FIR point of view, M87 is a passive object with a central radio source emitting synchrotron emission, without a substantial diffuse dust component.     \\begin{figure} \\centering \\includegraphics[width=0.485\\textwidth]{14555fg4.pdf} \\caption{A comparison between the synchrotron model image and the observed image at 250~$\\mu$m. The left panel shows the synchrotron image at the model resolution (8\\farcs5 FWHM, 2\\arcsec\\ pixels), the central panel shows the same model convolved to the SPIRE 250~$\\mu$m beam and pixel size (18\\farcs1 FWHM, 6\\arcsec\\ pixels). The right panel shows the observed SPIRE 250~$\\mu$m image. In all panels, the green lines are the contours of the synchrotron model at the model resolution.} \\label{H250.pdf} \\end{figure}      \\begin{acknowledgement} The National Radio Astronomy Observatory (NRAO) is operated by Associated Universities Inc, under cooperative agreement with the National Science Foundation. \\end{acknowledgement}"
    },
    "1005/1005.3745_arXiv.txt": {
        "abstract": " ",
        "introduction": "Since the last decade of the 20th century it is known that there are 340 observed extrasolar planetary systems with more than 400 planets and some of them have two or more planets (Schneider, October 2009). In many cases of multiplanetary systems, two planets are in mean motion resonance, (e.g.  HD 82943, GJ 876, HD128311, 55CnC$_{b,c}$). Some of these systems have large eccentricities and are evidently stable, since they exist in nature. Thus the role of the resonances in the evolution of a planetary system should play an important role either for its long term stability or its migration to the present position.  There are different approaches to the study of the long term dynamical stability of a resonant planetary system  and on the mechanisms that stabilize the system or generate chaotic motion and instability (e.g. Beaug\\'e et al, 2003; Ferraz-Mello et al, 2006;  Marzari et al, 2006; Callegari et al, 2006; Baluev, 2008; Gozdziewski et al, 2008; Beaug\\'e et al, 2008; Schward et al, 2009).  In these papers different methods have been applied, as the averaging method, direct numerical integrations of orbits, or various numerical methods which provide indicators for the exponential growth of nearby orbits. In this way the regions where stable motion exists have been detected, in the orbital elements space. Another method to study the dynamics of a planetary system is to find the basic families of periodic orbits and their stability. It is known that the periodic orbits play a dominant role in understanding the dynamics of a system, because they determine critically the structure of the phase space. In this way, we can detect the regions where stable librations could exist (Hadjidemetriou, 2006; Voyatzis, 2008). These are the regions close to the stable periodic orbits, where a real planetary system could be trapped. The periodic orbits with nonzero eccentricities of the planets, either stable or unstable, are associated with mean motion resonances between the planets. Also, work on these lines are by Hadjidemetriou (2002), Psychoyos and Hadjidemetriou (2005), Voyatzis and Hadjidemetriou (2005, 2006). A review on these issues is given in Hadjidemetriou (2008).  It is believed that the existence of a planetary system in a region of long term stability is the result of a dynamical process that forces the planets to migrate after their formation (see e.g. Tsiganis et al, 2005; Morbidelli et al, 2007).  Friction, due to the interaction between the newly formed planets and the protoplanetary nebula, before this latter is dissolved, is the main source of non conservative forces that generate a {\\it migration} of the planets. Also tidal friction due to the body deformation also produces non conservative forces that affects both the orbit and the spin evolution (Ferraz-Mello 2008). Non conservative forces may explain the large eccentricities that are observed in some planetary systems or the very small planetary distances from the sun.  The introduction of non conservative forces destroy the symplectic structure of a purely gravitational system and the phase space volume is not conserved. Consequently the overall dynamics and the topology of the phase space change critically. A particular review and relative literature can be found in Contopoulos (2002).  Numerical simulations using hydrodynamical models (e.g. Nelson and Papaloizou, 2003a,b) or dissipative restricted three body models (Beaug\\'e and Ferraz-Mello, 1993; Gomes, 1995; Haghighipour, 1999)  showed that planetary evolution leads to a trapping of planets in a mean motion resonance. Such evolution follows  paths of resonant stable configurations (Lee and Peal 2002; Papaloizou 2003; Lee, 2004; Sandor and Kley, 2006). Such paths are proved to be characteristic curves of resonant corotations (Ferraz-Mello et al, 2003; Beaug\\'e et al 2006; Zhou et al, 2008), which actually coinside with families of periodic orbits in a rotating frame (Voyatzis and Hadjidemetriou 2005, 2006). The present work aims to study the close association of a possible evolution of a planetary system under dissipative forces with the families of periodic orbits of the conservative model. As we shall see in the following, the families of periodic orbits play a crucial role on the evolution of the system, in phase space, under any dissipation law and the knowledge of the families of periodic orbits is necessary in understanding the migration process.  We consider, additionally to the gravitational forces, the action of the Stokes non conservative force \\begin{equation} \\vec{R}=-10^{-\\nu}(\\vec{v}-\\vec{v_c}),\\;\\;\\nu>0 \\label{diss1} \\end{equation} where the velocity $\\vec{v_c}$ is the Keplerian circular velocity at the distance $r$ from the Sun, given by $$ \\vec v_c=\\sqrt{\\frac {Gm_0}{r}}\\:\\vec{e_\\theta}. $$ The vector $\\vec e_\\theta$ is the unit vector normal to the radius vector $\\vec r$ and $m_0$ is the mass of the sun. The Stokes dissipation law (\\ref{diss1}) implies that the nebula that generates the dissipation rotates differentially with a Keplerian circular velocity $\\vec v_c$ at any distance $r$. We also consider a variation of the Stokes force given by \\begin{equation} \\vec{R}=-10^{-\\nu}(\\vec{v}-\\vec{v_c}/r^{1/2}),\\;\\; \\nu>0. \\label{diss2} \\end{equation} This law differs from the law given in Equation (\\ref{diss1}), in that the circular velocity is multiplied by $1/r^{1/2}$. This means that the  velocity of rotation of the protoplanetary nebula varies with the distance from the sun. It is larger than $\\vec v_c$ close to the sun and goes slowly to zero as $r\\rightarrow \\infty$. We remark that $\\vec{R}$ is applied to both planets and not only to the outer planet as e.g. in the studies of Lee and Peal (2002), Beaug\\'e et al (2006).  The model we use is the {\\it general three body problem}, for planar motion, with the sun, $S$, and the two planets, $P_1$ (inner) and $P_2$ (outer), as the three bodies of mass $m_0$, $m_1$ and $m_2$, respectively, normalized such that $m_0+m_1+m_2=1$ . We assume that the center of mass of the whole system is at rest with respect to an {\\it inertial} frame $XOY$. We have four degrees of freedom, for planar motion, with generalized variables $X_1,\\:Y_1,\\:X_2,\\:Y_2$. We define next, a {\\it rotating} frame of reference $xOy$, whose $x$-axis is the line $S-P_1$, with origin at the center of mass of these two bodies, where $S$ is the Sun and $P_1$ the inner planet. This is a non uniformly rotating frame, where the planet $P_1$ moves  on the $x$-axis and the planet $P_2$ moves in the $xy$ plane. As generalized coordinates in this rotating frame we can select $x_1,\\:x_2,\\:y_2$ and $\\theta$, where $\\theta$ is the angle between the $Ox$ axis and a fixed direction in the inertial frame. It turns out that the angle $\\theta$ is ignorable, so we have three degrees of freedom and the study of the system is restricted in the rotating frame $xOy$, with variables $x_1,\\:x_2,\\:y_2$ (Hadjidemetriou, 1975).  The evolution of the planetary system can be studied by computing the Poincar\\'e map on a surface of section. By this method we reduce the dimensions of the phase space, without losing the generality of the problem. In the present study, we consider the surface of section \\[y_2=0 \\:\\:(\\dot y_2>0),\\] in the rotating frame $xOy$.  It can be shown (see e.g. Voyatzis and Hadjidemetriou 2005, 2006) that families of symmetric or asymmetric periodic orbits exist in the {\\it rotating} frame $xOy$. In particular, in a symmetric periodic orbit the planet $P_2$ intersects perpendicularly the $x$-axis while the planet $P_1$ is temporarily at rest on the $x$-axis. Consequently, the non zero initial conditions of a symmetric periodic orbit are $x_{10}$, $x_{20}$, $\\dot y_{20}$, and a family of symmetric periodic orbits is represented by a smooth curve in the space $\\Pi_3=\\{x_{10},x_{20},\\dot y_{20}\\}$ of initial conditions. An asymmetric periodic orbit is represented by a smooth curve in a five dimensional space of initial conditions e.g. $\\Pi_5=\\{x_{10}, x_{20},\\dot{x}_{10},\\dot x_{20}, \\dot y_{20}\\}$. In the following we shall present a family by its projection on the coordinate plane $x_{10}x_{20}$ or in the space $e_1e_2$ of the osculating eccentricities at $t=0$. When we plot symmetric periodic orbits in the space $e_1e_2$ we use the convention $e_i>0$ or $e_i<0$ when the planet is at aphelion or at perihelion, respectively.  In the conservative system the stable fixed points of the Poincar\\'e map correspond to periodic orbits in the rotating frame. If in the equations of motion we consider dissipative forces, in addition to the gravitational forces, the energy integral does not exist. The stable fixed points are now stable limit cycles, to which the system is attracted in some cases. Chaotic attractors may also exist and correspond to a set of points (as $t\\rightarrow \\infty$), which are irregularly distributed in the section. ",
        "conclusions": "In this work we studied the evolution of a planetary system with two planets moving in the same plane, under the action of dissipative forces. The main study was made for the Stokes like force given by Eq. (\\ref{diss1}), for different values of $\\nu$. This force implies that the two planets move inside a protoplanetary nebula that rotates with a differential Keplerian circular velocity $\\vec v_c$. We also considered a variation of this dissipative force given by eq. (\\ref{diss2}) and, for the 3/2 resonance, also the force $\\vec{R}=-10^{-\\nu}(\\vec{v}-\\vec{v_c}/r^{3/2})$. In these last cases the rotational velocity of the protoplanetary nebula deviates from the circular Keplerian velocity, being larger close to the sun and tending to zero at large distances.  An important result that was obtained in the present study is that the families of periodic orbits play a crucial role in the evolution of the system: These families, both stable or unstable, determine the structure of the phase space. It was shown that the {\\it stable} families provide the routes along which the system evolves, while the unstable families define the regions that the system avoids in its evolution. {\\em The evolution may be different for different dissipation laws, but in all cases it is the families of periodic orbits that guide the system in its evolution, either to smaller or to larger eccentricities}. The particular dissipation law determines how the system moves along the families and its final trapping to a resonance.  We studied the evolution of a planetary system considering two different cases.  Case 1: The planetary system starts far from a periodic motion (either resonant or non resonant).\\\\ Case 2: The planetary system starts close to a resonant periodic motion. In particular, we studied the 2/1 and 3/2 resonances.  The evolution depends on the particular dissipative force.  \\underline{Case 1}:\\\\ We considered first the Stokes like law of eq. (\\ref{diss1}). We found that in all cases the system is attracted to an orbit on the {\\it circular family of periodic orbits}, following an irregular route with almost the same value of the planetary frequencies $n_1/n_2$ as the starting value. After this, the system evolves {\\it along} the circular family, in a smooth way, to the direction where $n_1/n_2$ decreases, until it meets a significant resonance, for example 2/1 or 3/1. From that point on, the system is trapped in a chaotic attractor or in a resonant orbit, respectively, with very small eccentricities (see figures \\ref{e12-28}, \\ref{e12-35}). The semimajor axes increase slowly and then are stabilized at the values corresponding to the resonant attractor. This evolution is the same for a whole region in the eccentricity space, which is in fact the {\\it basin of attraction} to the above mentioned chaotic attractor (Fig. \\ref{e12-35-all}) for the 3/1 resonance or to a circular resonant orbit for the 2/1 resonance. Thus we get a mechanism of trapping to a resonance with very small eccentricities.  A different evolution appears for the dissipation law (\\ref{diss2}). The system moves towards a 3/1 resonant orbit on the circular family and then follows the {\\it stable} family of elliptic 3/1 resonant periodic orbits that bifurcates from the resonant circular orbit. It evolves along this family, with {\\it increasing} eccentricities, and if the family becomes unstable, it changes route and follows the stable asymmetric family that bifurcates from this critical point. This process provides a mechanism of trapping to an asymmetric 3/1 resonant periodic orbit with large eccentricities.  \\underline{Case 2}:\\\\ We start with the Stokes like eq. (\\ref{diss1}), as in case 1. We study first the evolution close to the 2/1 resonance. In all cases, the evolution follows the {\\it stable} parts of the family of 2/1 resonant periodic orbits, with {\\it decreasing} eccentricities (see figures \\ref{fam21-1}b, \\ref{fam21-3}b for the 2/1 resonance and Fig. \\ref{e12-32}a for the 3/2 resonance). The semimajor axes are stabilized to fixed values, after a slow decrease. If, however, there is a chaotic region along the family, then the system develops chaotic motion (unless the dissipation process stops earlier) with decreasing eccentricities (Fig. \\ref{fam21-2}). This evolution is the same either for a starting orbit on the 2/1 family or far from it.  The evolution is different for the dissipative force given by eq. (\\ref{diss2}). Depending on the starting point, the eccentricities decrease or increase and the system is finally trapped to a resonant periodic orbit with nonzero eccentricities (see Fig.\\ref{e12-2-21} for the 2/1 resonance and Fig. \\ref{e12-32}b for the 3/2 resonance).  Finally, a different evolution appears for the dissipative force $\\vec{R}=-10^{-\\nu}(\\vec{v}-\\vec{v_c}/r^{3/2})$. Contrary to the previous two cases, the eccentricities in this case {\\it increase}, following the stable parts of the resonant family, Chaotic motion may appear if an unstable region is met on the family (see Fig.\\ref{e12-32}c for the 3/2 resonance).  We see that the final outcome, whether the system is trapped to small or large eccentricities, depends critically on the particular dissipation force, and in particular on how much the rotational velocity of the protoplanetary nebula deviates from the Keplerian circular velocity. In almost all cases the system is trapped to a resonant orbit."
    },
    "1005/1005.4353_arXiv.txt": {
        "abstract": "We present the optical spectra of a sample of 34 SWIRE-CDFS sources observed with EFOSC2 on the ESO 3.6m Telescope. We have used the spectra and spectroscopic redshifts to validate our photometric redshift codes and SED template fitting methods. 12 of our sources are Infrared Luminous Galaxies. Of these, five belong to the class of ULIRGs and one to the class of HLIRGs with evidence of both an AGN and starburst component contributing to their extreme infrared luminosity for 3, starburst contributing for 1 and AGN contributing for 2 of them. ",
        "introduction": "In recent years, extragalactic surveys have revolutionized our view of the high-redshift Universe, showing that the redshift range $0.5<z<3.0$ witnessed an extraordinary transformation in the demographics of starburst and AGN activity in galaxies. The factors that drove these changes are however unclear. Consequently, the properties of large samples of sources are the current driving force behind modern extragalactic surveys (e.g. Lonsdale et al. 2003; Eales et al. 2010; Oliver et al. 2010). However, either most of the sources in these surveys are too faint for spectroscopic observations or it is not cost and time effective to acquire large spectroscopic redshift samples using the current facilities. Photometric techniques provide an excellent tool for estimating redshifts, essential for statistical studies of the galaxies' properties and evolution with only a fraction of the observing time required by spectroscopic surveys. These techniques are faster and can estimate redshifts beyond the spectroscopic limit with the main handicap of possessing larger errors than spectroscopic techniques. In order to minimize these statistical errors: i) more bands are needed to be included in the photometric computation ii) optical spectroscopic redshifts have to be used to calibrate and validate the photometric redshift codes. \\\\ In this paper we use the EFOSC2 on the 3.6m ESO telescope to obtain optical spectroscopy of galaxies selected from the SWIRE-CDFS in order to validate our photometric redshift techniques across a wide redshift range, thus allowing more accurate determinations for the whole SWIRE catalogue. In addition, we present the optical spectra and MIR properties of a sample of LIRGs detected as part of our follow-up. Infrared luminous objects at $z>0.5$ are of great importance since they lie at the proposed peak of global starburst activity. They may well be related to the famous sub-mm galaxy population (eg. Chapman et al. 2003) and be connected to the origin of the Cosmic Infrared Background (Puget et al. 1996). Many questions remain about this population, including the role of AGN in their prodigious luminosities, what their descendants are in the local universe, and what their broader role is in galaxy formation and evolution. Some have claimed, for example, that such far-IR luminous objects can only be incorporated into our general understanding of galaxy formation if they have stellar initial mass functions (IMFs) skewed to high masses (Baugh et al. 2005).   \\begin{figure*} \\centering \\begin{center} \\includegraphics[scale=0.43]{redshift.ps} \\includegraphics[scale=0.43]{Rmag.ps} \\end{center} \\caption{\\textbf{Left} : The redshift distribution of the 34 extragalactic sources of the spectroscopic sample. \\textbf{Rigth} : R-band distribution for the same spectroscopic sample.} \\label{fig:Redshift_distribution} \\end{figure*} ",
        "conclusions": "We present the optical spectra of 34 sources  within SWIRE-CDFS, observed using EFOSC-2 in MOS mode on the ESO 3.6m telescope. We have compared the observed spectroscopic redshifts values to those estimated using our photoz methods and the agreement is very good especially in the case of galaxies. With respect to QSOs' photoz agreement and taking into consideration the problems photoz methods face when dealing with quasars, our sample is too small to extract safe results. Among our sources we found 6 LIRGs, 5 ULIRGs and 1 HLIRG. All H/ULIRGs are broad line objects at z$>$1.0, in excellent agreement with our SED template fitting method which fitted all of them with a QSO optical template. All the broad line objects in our sample are fitted with a QSO optical template and show evidence of a strong dust torus component in the MIR in excellent agreement with MIR color-color plots. In the case of narrow emission line objects, based on emission diagnostic diagrams for the 15 sources with the necessary lines, we found 6 pure star-forming galaxies, 3 Seyfert galaxies, 6 composite objects of which 3 appear as star-forming objects and 3 as LINERs based on the [\\textit{SII}]/\\textit{H$\\alpha$} vs [\\textit{OIII}]/\\textit{H$\\beta$} diagram."
    },
    "1005/1005.2952_arXiv.txt": {
        "abstract": "{With a luminosity $>10^5$ \\Lsun\\ and a mass-loss rate of $\\sim 2. 10^{-4}$ \\Msun\\,yr$^{-1}$, the red supergiant VY CMa truly is a spectacular object. Because of its extreme evolutionary state, it could explode as supernova any time. Studying its circumstellar material, into which the supernova blast will run, provides interesting constraints on supernova explosions and on the rich chemistry taking place in such complex circumstellar envelopes. We have obtained spectroscopy of VY\\,CMa over the full wavelength range offered by the PACS and SPIRE instruments of Herschel, i.e. 55 -- 672 micron. The observations show the spectral fingerprints of more than 900 spectral lines, of which more than half belong to water. In total, we have identified 13 different molecules and some of their isotopologues. A first analysis shows that water is abundantly present, with an ortho-to-para ratio as low as $\\sim$1.3:1, and that chemical non-equilibrium processes determine the abundance fractions in the inner envelope.}{}{}{}{} ",
        "introduction": "The oxygen-rich red supergiant VY\\,CMa is one of the most extreme objects of its kind. Its large mass-loss rate results in a very dense circumstellar envelope and an enormous infrared excess \\citep[][]{2001ApJ...557..844H}. Infrared observations are invaluable for in-depth studies of the complex chemical and dynamical processes at work in its circumstellar envelope. By this means, VY\\,CMa has been shown to harbour a rich, water-dominated chemistry \\citep[e.g.][]{1999ApJ...517L.147N, Polehampton2010A&A...510A..80P}. The unprecedented wavelength coverage and spectral resolution of the Herschel instruments will undoubtedly push our understanding of the mass-loss related phenomena a step further.%  In this context, VY\\,CMa has been observed as part of the Herschel   Guaranteed Time Key Programme ``Mass-loss of Evolved StarS'', MESS (Groenewegen et al., {\\it in prep.}). The spectra discussed hereafter have been obtained with the PACS and SPIRE instruments of Herschel \\citep{Pilbratt2010}. The instruments and their in-orbit performance are described in \\citet{Poglitsch2010} and \\citet{Griffin2010}. The SPIRE calibration methods and accuracy are outlined by \\citet{Swinyard2010}.  \\begin{figure*} \\centering \\includegraphics[width=\\textwidth,height=16.3cm]{14641fig1.eps} \\caption{The complete PACS + SPIRE spectrum of VY~CMa. The various spectral bands are presented in order of increasing wavelength. For the sake of clarity, PACS band B2A is not included, as it matches the wavelength range of PACS-B3A. The ISO-LWS spectrum is overplotted for comparison \\citep{Polehampton2010A&A...510A..80P}. For PACS, the spectrum of the central spaxel is displayed, after application of the point source correction (version 2.0). The strongest spectral features of the main molecules are indicated.} \\label{FigFullSpectrum} \\end{figure*} ",
        "conclusions": "Through the detection of more than 900 emission lines of various chemical species in VY\\,CMa, PACS and SPIRE have shown their diagnostic strength to unravel the chemically complex mass-loss processes at work in evolved stars. Specifically, we have derived the HCN, SiO and water fractional abundances of the circumstellar material around VY\\,CMa, which also shows a suprisingly low ortho-to-para water ratio, close to 1.  The origin of the latter is not yet understood, but the Herschel observations of many evolved stars in the MESS programme might shed light on this topic."
    },
    "1005/1005.2213_arXiv.txt": {
        "abstract": "{In this first paper on the results of our \\emph {Herschel} PACS survey of local Ultraluminous Infrared Galaxies (ULIRGs), as part of our \\emph {SHINING} survey of local galaxies, we present far-infrared spectroscopy of Mrk 231, the most luminous of the local ULIRGs, and a type 1 broad absorption line AGN.  For the first time in a ULIRG, all observed far-infrared fine-structure lines in the PACS range were detected and \\emph{all} were found to be deficient relative to the far infrared luminosity by 1 -- 2 orders of magnitude compared with lower luminosity galaxies.  The deficits are similar to those for the mid-infrared lines, with the most deficient lines showing high ionization potentials.  Aged starbursts may account for part of the deficits, but partial covering of the highest excitation AGN powered regions may explain the remaining line deficits.  A massive molecular outflow, discovered in OH and $^{18}$OH, showing outflow velocities out to at least 1400 km sec$^{-1}$, is a unique signature of the clearing out of the molecular disk that formed by dissipative collapse during the merger.  The outflow is characterized by extremely high ratios of $^{18}$O / $^{16}$O  suggestive of interstellar medium processing by advanced starbursts. } ",
        "introduction": "ULIRGs (L$_{IR}$ $\\geq 10^{12} L_{\\odot}$), which in the local universe are known to be mergers of gas rich galaxies, are thought to play a critical role in galaxy evolution.  It is therefore crucial to understand the conditions, dynamics, chemistry and energetics of this stage of evolution, in which dissipative collapse accompanies the transformation of gas rich galaxies into ellipticals (Lonsdale et al. 2006). As part of the \\emph{SHINING} Key Project on local galaxies, we are carrying out a far-infrared (FIR) spectroscopic survey of all 21 ULIRGs in the Revised Bright Galaxy Survey (Sanders et al. 2003).  Previous work on ULIRGs with the Infrared Space Observatory (ISO) Long Wavelength Spectrometer (LWS) found deficits in far-infrared atomic and ionized fine-structure line emission relative to their FIR luminosities (Luhman et al. \\cite{Luhman98}) accompanied by prominent molecular absorption in excited transitions of molecules such as OH and H$_2$O rarely seen in other galaxies (Fischer et al. \\cite{Fischer99}).   The deficits, often based only on a few non-detections, were postulated to result from the high ratios of UV radiation density to particle density in the nuclei of these galaxies (Malhotra et al. 2001, Luhman et al. 2003, Abel et al. 2009), high gas density (e.g. Negishi et al. 2001)  and high FIR opacity and/or high luminosity-to-mass ratio (Gonz\\'{a}lez-Alfonso et al. 2004, 2008, hereafter GA04, GA08), but because of the paucity of diagnostic line detections this puzzle remained unresolved.  GA04 and GA08 showed that the high excitation molecular absorptions in both Arp 220 and Mrk 231 are radiatively pumped, consistent with high radiation density to particle density. They identified far-infrared absorption by OH, H$_2$O, NH$_3$, NH, and $^{18}$OH, but the identifications of the latter species were uncertain due to the low resolution of the LWS.  The PACS ULIRG spectroscopic survey is designed to kinematically identify the ionized, atomic and molecular regions and to study the conditions of the interstellar medium, thereby illuminating the nature of this important evolutionary phase.  In this first paper we present our initial results on Mrk 231, the most luminous of the local ULIRGs (L(8--1000 $\\mu$m) = 3.2 $\\times$ 10$^{12}$ L$_{\\odot}$) and a type 1, low-ionization broad absorption line (LoBAL) active galactic nucleus (AGN) at an adopted distance of 172 Mpc (z=0.04217).  Its central quasar is covered by a semi-transparent dusty shroud producing about 3.1 magnitudes of extinction at 4400 $\\AA$ (Reynolds et al. 2009) and is at the center of a rotating, nearly face-on molecular disk (Downes \\& Solomon 1998).  Based primarily on \\emph {Spitzer} results, Veilleux et al. (2009) estimate that the average AGN contribution to the bolometric luminosity in ULIRGs is 35 -- 40\\% and that for Mrk 231 the AGN contribution is $\\sim$ 70\\% by most estimation techniques.  The contribution of an advanced 120 -- 250 Myr nuclear starburst is estimated at 25 -- 40\\% based on near infrared observations of Mrk 231 by Davies et al. (2007).  \\begin{table*} \\caption{Spectroscopic measurements for Mrk 231} \\label{Table1} \\centering \\begin{tabular}{r ccccc}     % \\hline\\hline Transition\\tablefootmark{a,}\\tablefootmark{b} & Line flux  & Line flux error\\tablefootmark{c} & Equivalent Width & Measured FWHM & Inferred FWHM\\tablefootmark{d}\\\\ & 10$^{-17}$ W m$^{-2}$\t & 10$^{-17}$ W m$^{-2}$\t& 10$^{-3}$ $\\mu$m\t& km s$^{-1}$\t& km s$^{-1}$ \\\\ \\hline  [N III] 57.3 $\\mu$m & 2.8  & 0.6 & -1.0 & 203 & 177\\\\  [O I] 63.2 $\\mu$m & 36.0 & 2.6 & -17.0 & 232 & 218\\\\  [O III] 88.4 $\\mu$m & 4.1 & 0.6 & -3.8 & 311 & 289\\\\  [N II] 121.9 $\\mu$m & 4.1 & 0.5 & -10.4 & 389 & 266\\\\  [O I] 145.5 $\\mu$m & 3.2 & 0.4 & -17.9 & 324 & 208\\\\  [C II] 157.7 $\\mu$m & 38.3 & 1.3 & -296.0 & 337 & 247\\\\  [N II] 205.2 $\\mu$m\\tablefootmark{e} & 2.6 & 0.2 & -61.0 & 524 & 242\\\\  OH 79.1, 79.2 $\\mu$m & -22.8 & 1.0 & 15.3 & -- & --\\\\ $ $ & 3.2 & 0.6 & -2.1 & -- & --\\\\  OH 119.2, 119.4 $\\mu$m & - 47.6 & 0.3 & 107.4 & -- & --\\\\ $ $ & 9.4 & 0.3 & -21.6 & -- & --\\\\  $^{18}$OH 120.0, 120.2 $\\mu$m & -14.7 & 0.2 & 34.3 & -- & --\\\\ $ $ & 1.7 & 0.2 & -4.0 & -- & --\\\\  H$_2$O 78.7 $\\mu$m & -18.0 & 0.1 & 12.2 & 444 & 422\\\\  HF $\\diagup$ H$_2$O 121.7 $\\mu$m & -4.7 & 0.4 & 11.8 & 560 & 483\\\\  \\hline \\end{tabular}\\\\ \\tablefoottext{a}{For doublets, fluxes listed are for the blended pair.} \\tablefoottext{b}{For transitions with both absorption and emission, the absorption component is listed first.} \\tablefoottext{c}{Errors listed are the statistical rms uncertainties;  we allow 25\\% calibration uncertainties (see text).} \\tablefoottext{d}{Inferred velocity width is based on the instrument resolution assuming a Gaussian profile;  FWHMs are not listed for doublets.} \\tablefoottext{e}{{Herschel} SPIRE data (V10).} \\end{table*} ",
        "conclusions": "\\subsection{The fine-structure lines}  The continuum normalized fine-structure line spectra are displayed in Fig.~\\ref{FineStructureLines} as a function of velocity, where Mrk 231's redshift of 0.04217 is defined as 0 km s$^{-1}$ and the red horizontal bars represent the instrumental FWHMs.  The FWHM velocities observed in these lines are in all cases larger than the instrumental resolution and we do not detect self-absorption.  The average of the derived intrinsic FWHMs for the fine-structure lines is 235 km s$^{-1}$, similar to the FWHM velocity widths of 167 km s$^{-1}$ and  270 km s$^{-1}$ measured by Sanders et al. (1991) and Tacconi et al. (2002) from CO(1--0) and near-infrared stellar CO bandhead studies, respectively.  Thus the fine-structure line emission is probably associated with the central star forming disk.  In  the [CII] line, wings not characteristic of the instrumental profile are present, out to $-1000$ km s$^{-1}$ for the blue wing, suggesting that some C$^+$ is associated with the outflow discussed in the next subsection.  The [NII]205 profile also suggests a wing out to $-700$ km s$^{-1}$, consistent with the detection of optical [NII] lines by Rupke et al. (2005).  We  estimate that the absorption by HF or H$_{2}$O observed in the blue wing of the [NII] 122 $\\mu$m line does not affect the [NII] line flux by more than 50\\%, since the FWHM of this transition is similar to those of the other fine-structure lines.  The observed [NII]205 to [NII]122 ratio of 0.6 $\\pm$ 0.4 translates into a density range for N$^{+}$ of  log n$_{e}$ = $1.4^ {+0.8}_{-0.4}$ (Rubin et al. 1993).  Photodissociation region (PDR) modeling of [CII], [OI], and FIR fluxes yields UV radiation field and density $G_{0}$ $\\approx$ 6 $\\times$ 10$^{3}$ and $n$ $\\approx$ 500 cm$^{-3}$, at the extreme range of average radiation density per particle density seen in extragalactic nuclei (Sturm et al. 2010, Fig. 4).  The fine-structure line strengths relative to the far-infrared luminosities are weak compared with normal HII galaxies and AGN.  In Fig.~\\ref{LineDeficits}, we plot as a function of wavelength (top) and ionization potential (bottom) the fine-structure line to far-infrared flux ratio (FIR, 42 -- 122 $\\mu$m) in Mrk 231 relative to the median value for subsamples of HII galaxies (black symbols) and AGN (red symbols) with luminosities 1 x 10$^9$ L$_\\odot$ $\\leq$ L$_{IR}$ $\\leq$ 5 x 10$^{11}$ L$_\\odot$ and major isophotal diameters D$_{25}$ $\\leq$ 250 arcsecs compiled from the literature and our \\emph {SHINING} sample (Graci\\'{a}-Carpio et al., in prep.).  There are no [NII] 205 $\\mu$m line measurements for the comparison samples, so this line is not included in Fig.~\\ref{LineDeficits}.  The observed deficits are severe:  typically an order of magnitude for the neutral and low ionization species and up to two orders of magnitude for the higher ionization lines.  There is no correlation between the deficit and wavelength, and only weak correlation with critical density, so \\emph{differential} extinction (with wavelength) does not appear to play a significant role, and density does not play a major role.  We find a strong inverse correlation (correlation coefficient, $-0.78$) with ionization potential when compared with the AGN sample.  A similar inverse correlation is found with the starburst comparison sample only if the [NeIII]15 deficit is treated as an outlier.  Low excitation is consistent with the estimation by Davies et al. (2007) that a 120 -- 250 Myr starburst supplies 25 -- 40 \\% of the bolometric luminosity of Mrk 231.  However such an aging starburst can create  a deficit of 40\\% at most.  The high relative strength of the [OI] 145 $\\mu$m line, with its high excitation lower level, is suggestive of increased radiation density per particle density, as parameterized by the ionization parameter $U$ (Luhman et al. 2003; Abel et al. 2009).  The high ratio [OI] 145 / [CII] 158 $\\approx$ 1/11 lies between the predicted line ratio for starburst and AGN at the far-infrared F(60)/F(100) color of Mrk 231 with $U$ $\\approx$ 10$^{-2}$ (high $U$) in the dust-bounded models of Abel et al. (2009).  The strong inverse correlation relative to the AGN comparison sample in Fig.~\\ref{LineDeficits}  is consistent with partial covering, higher for high ionization lines, of the line emitting region, due for example to a face-on molecular torus, with optically thick clumps.  \\begin{figure} \\centering \\includegraphics[width=6.2cm]{fig3.pdf} \\caption{P-Cygni profiles (solid blue, black histograms) of the OH 119 $\\mu$m and $^{18}$OH 120 $\\mu$m partially resolved doublets (\\emph{top}) and the OH 79 $\\mu$m unresolved doublet (\\emph{bottom}) in Mrk 231 are compared with our modeled profiles (solid red curves, see text).  Absorption in the H$_2$O 4$_{23}$ -- 3$_{12}$ transition with a possible blue-shifted wing is modeled in GA10.  The zero velocity rest-frame wavelengths are marked with black (OH), blue ($^{18}$OH), and magenta (H$_2$O) arrows and the instrumental FWHM is indicated in each panel with horizontal red bars. } \\label{MolecularLines} \\end{figure}  \\subsection{The molecular outflow}  Figure~\\ref{MolecularLines} displays the observed 119 -- 120 $\\mu$m OH / $^{18}$OH and 79 $\\mu$m H$_2$O / OH line profiles in Mrk 231. The OH lines show spectacular P-Cygni profiles in all three ground-state doublets, with extremely broad blue-shifted absorption as far out as 1400 km s$^{-1}$ in the OH 119 $\\mu$m line.  Both the high \\emph {molecular} outflow velocities and the relative strengths of the OH and $^{18}$OH are unprecedented.  The excited H$_{2}$O 78.7 $\\mu$m and HF/H$_{2}$O 121.7 $\\mu$m lines (see Fig.~\\ref{FineStructureLines}) appear to have blue-shifted wings, indicating that H$_2$O and HF may also be involved in the outflow.  Recently a molecular outflow was detected by Sakamoto et al. (2009) in submillimeter lines of HCO$^+$ and CO toward the nuclei of Arp220 with an outflow velocity of $\\sim100$ km s$^{-1}$, but this is modest compared with that of Mrk 231.  Our model of the observed profiles is shown in Fig.~\\ref{MolecularLines} and is based on the continuum components and radiative pumping described in Gonz\\'{a}lez-Alfonso et al. (2010, hereafter GA10).  It uses three continuum components, hot (radius = 23 pc, $T_{dust}$ = 400 - 150 K), warm (120 pc, 95 K), and extended (610 pc, 41 K), and three velocity components at $100$, $400$, and $600$ km s$^{-1}$ (the latter component with high turbulent velocity, $300$ km s$^{-1}$).  We have conservatively assumed a screen approach for OH, with the outflow located just around the warm continuum source, resulting in $N(\\mathrm{OH})=1.4\\times10^{17}$ cm$^{-2}$. The model indicates mechanical energy $\\gtrsim10^{56}$ erg (dominated by the high velocity component), mechanical luminosity $\\gtrsim1$\\% of the TIR luminosity, and an outflow mass of $\\gtrsim7\\times10^7$ M$_{\\odot}$.   Observations of high-lying OH lines would help better constrain the outflow location, energetics, and the apparent extremely low $^{16}$O / $^{18}$O ratio.  Figure~\\ref{OpticalDepths} (top) shows that the absorption at 120 $\\mu$m by $^{18}$OH at low outflow velocities ($>-250$ km s$^{-1}$) is deeper than the corresponding absorption in the cross-ladder 79 $\\mu$m OH doublet at the same velocity.   In this simple model we assume that the incomplete absorption in the optically thick OH 119 $\\mu$m doublet is  a result of the fact that OH at a given velocity is only covering a fraction of the continuum (Fig.~\\ref{OpticalDepths}, middle) and we ignore any differential re-emission of the $^{18}$OH and 79 $\\mu$m OH lines.   We infer that at velocities $>-250$ km s$^{-1}$ $\\tau_v(\\mathrm{^{18}OH \\, 120 \\mu m}) > \\tau_v(\\mathrm{OH \\, 79 \\mu m})$ (Fig.~\\ref{OpticalDepths}, bottom).  From the ratio of the Einstein B coefficients, we find that the opacity in the 79 $\\mu$m OH line is 40 times lower than in the 119 $\\mu$m OH line and thus that $N(\\mathrm{OH})/N(\\mathrm{^{18}OH}) < 40$.  These results are not critically dependent on the uncertainties in the covering factor due to re-emission (or thermal emission) in the 119 $\\mu$m OH line, which would affect both the opacities of the $^{18}$OH and 79 $\\mu$m OH lines but not their relationship.  We note that the effects of filling in by cooler dust would weaken the 120 $\\mu$m absorption more than the 79 $\\mu$m absorption.  The rarer isotopologue $^{17}$OH would appear between the two lines at 119 and 120 $\\mu$m, but it is not possible to put a limit on its strength due to the broad width of the detected features.  We note that CH$^+$, HCl, H$^{37}$Cl, H$^{79}$Br, and H$^{81}$Br have transitions at $\\approx120$ $\\mu$m, but HCl and HBr and their isotopologues are not expected to significantly contaminate the observed absorption due to the high lower level energies of their $120$ $\\mu$m transitions ($180-240$ K). On the other hand, CH$^+$(3-2) at 119.86 $\\mu$m could significantly contaminate the blue-shifted component of the $^{18}$OH doublet, given that CH$^+$(1-0) is detected in Mrk 231 (V10).  We estimate that contamination by CH$^+$ could decrease the implied $^{18}$OH abundance by a factor of $1.5-2$ and the $^{16}$O / $^{18}$O ratio could be increased to $\\sim$ $30-40$, still significantly below the solar system isotopic ratio of 500, the galactic center ratio of 250, the starburst galaxy ratio of $\\sim$ 150, and very close to the enhancement to $\\sim$ 50 that is theoretically predicted with a top heavy initial mass function in an advanced starburst stage (Henkel \\& Mauersberger 1993).  If our simple model is correct, the high $^{18}$O / $^{16}$O ratio seen here provides stringent constraints for stellar nucleosynthesis models.  Rupke et al. (2005) find that although both AGN and starburst galaxies have massive outflows traced  by broad interstellar Na I D absorption lines, the high velocity ($\\Delta$v$\\geq$ 2000 km s$^{-1}$) outflows seen almost exclusively in type 1 ULIRGs including Mrk 231, are probably powered by the AGN.  This was also shown in models by Narayanan et al. (2008).  If it is also true for the massive and energetic moderate velocity outflow seen here for the first time in OH and $^{18}$OH, then an AGN driven outflow is in the process of dispersing into the intergalactic medium, highly processed  material produced by an advanced starburst.  The molecular outflow shows velocities similar to the kiloparsec scale outflow component of Mrk 231 seen by Rupke et al. and appears to be a unique signature of the clearing out of the molecular disk that formed with the dissipative collapse during the merger.     \\begin{figure} \\centering \\includegraphics[width=6.8cm]{fig4.pdf} \\caption{Normalized velocity profiles for the blue-shifted OH and $^{18}$OH absorption (\\emph {top}).  The velocity scale is given relative to the bluest component of each OH doublet.  The covering factor (\\emph {middle}), assumed to be the same for OH and $^{18}$OH and derived from the profile of the optically thick OH 119 $\\mu$m line, is used to calculate the opacities in the 79 $\\mu$m and 120 $\\mu$m lines (\\emph {bottom}, see text). } \\label{OpticalDepths} \\end{figure}"
    },
    "1005/1005.5605_arXiv.txt": {
        "abstract": "In this paper we complete the program of the noncomutative geometry inspired black holes, providing the richest possible solution, endowed with mass, charge and angular momentum. After providing a prescription for employing the Newman-Janis algorithm in the case of nonvanishing stress tensors, we find regular axisymmetric charged black holes in the presence of a minimal length. We study also the new thermodynamics and we determine the corresponding higher-dimensional solutions. As a conclusion we make some consideration about possible applications. ",
        "introduction": "When we talk about curvature singularities in general relativity, we often forget that they are just fictitious effects emerging where the classical description of the gravitational field breaks down. In other words, to obtain physically reliable scenarios in spacetime regions plagued by curvature singularities, one has just the only possibility of invoking some formulation of the quantum theory of gravitation. Along this line of reasoning, new black hole geometries have successfully been derived, by improving the classical singularities with quantum gravity mechanisms. For instance, loop quantum black holes (LQBHs) have been obtained starting from the prescriptions of  loop quantum qravity. Inside the event horizon, a quantized version of the Kantowski-Sachs spacetime has replaced the irregular Schwarzschild geometry \\cite{LQBHs}. Other examples are the noncommutative geometry inspired black holes (NCBHs), whose line element singularities are cured by the presence of a minimal length effectively induced by the noncommutative character of coordinate operators \\cite{review,Nicolini:2005vd,NCBHs,NCBHs2}. Even if LQBHs and NCBHs use different tools to tame the curvature singularity, they share many common features, which would suggest an universal character when quantum gravity effects are taken into account. For instance, under some additional hypotheses, they both admit a maximum temperature, showing a thermodynamically stable final phase of the Hawking evaporation, or they both can be dirty, namely admitting  $g_{00}$ metric components which do not coincide with the usual value $-g_{11}^{-1}$ \\cite{dirty}. While LQBHs are for now known only for neutral, nonrotating, static four-dimensional geometries, the program of NCBHs includes a richer variety of physical situations: indeed the charged solution \\cite{Ansoldi:2006vg}, the extradimensional neutral \\cite{rizzo} and charged solutions \\cite{extracharged} and the rotating one \\cite{kerrr} have been obtained and they all exhibit a  deSitter core/belt as a consequence of quantum  manifold fluctuations at the origin. The general strategy to derive NCBH solutions consists of prescribing an improved form of the energy-momentum tensor, which accounts for the noncommutative fluctuations of the manifold at the origin and being vanishing for distances larger with respect to the noncommutative geometry typical scale, i.e. the minimal length. Another requirement is the covariant conservation of the energy-momentum tensor, while energy conditions at the origin are generally violated, in agreement with the fact that we are invoking a non-standard, i.e. non-classical, kind of matter. The derivation of the correct form of the energy-momentum tensor requires great care, in particular, for the axisymmetric case. Indeed, thanks to some considerations about the Kerr-Schild decomposition and some physical requirements for the regularity of the manifold, the noncommutative inspired Kerr (NCKerr) solution has been derived starting from a spinning fluid-type energy-momentum tensor. On the other hand, even for the spherically symmetric case, the inclusion of the electromagnetic interaction, makes the determination of the energy-momentum tensor much more difficult. In particular, if we want to have a complete scenario for a regular black hole endowed with all the three externally observable classical  parameters, mass, electric charge, and angular momentum, we need to devise an alternative to the direct derivation of the energy-momentum tensor. In this paper, we propose a way to circumvent the aforementioned difficulties opening a new route on the basis of the Newman-Janis algorithm \\cite{NJA}. However this kind of algorithm works only for vacuum or Reissner-Nordstr\\\"{o}m geometries. As a result, we first provide a prescription to amend the Newman-Janis algorithm and extend its validity to the case of NCBHs. Then we show how this prescription lets us correctly obtain the recent regular spinning black hole, i.e. the NCKerr solution. Then we proceed by deriving the corresponding regular Kerr-Newman (NCKN) solution and the higher-dimensional version of NCKerr and NCKN metrics. We recall that there exist other attempts to determine regular spinning black hole solutions \\cite{otherkerrr}. However the strategy at the basis of these derivations has been always that of matching an outer geometry with a regular local geometry at the origin. Conversely, by means of the Newman-Janis mechanism we shall present unique solutions which interpolate the local regular behavior at the origin and the conventional rotating geometries at large distances. In addition, against the existing literature, our solutions have an intrinsic quantum gravitational character at the origin, since the Newman-Janis algorithm maintains the original features of the corresponding nonrotating NCBHs.  The paper is organized as follows.  In Sec. \\ref{NJA}, we present the Newman-Janis algorithm which we will use in the following sections to derive the NCKerr solution (Sec. \\ref{kerrr}), the NCKN (Sec. \\ref{kerrrn}) and the higher-dimensional NCKerr and NCKN solutions (Sec. \\ref{extrakerrr}). Throughout the sections, discussions are devoted both to regularity of the geometry and the behavior of the Hawking temperature. In Sec. \\ref{concl}, we draw the conclusions. ",
        "conclusions": "\\label{concl} In this paper we have derived regular rotating charged and higher-dimensional black hole solution in the presence of a noncommutative geometry induced minimal length \\cite{NCQFT,NCQFT2}. To get this result, we reviewed the Newman-Janis algorithm in order to include a wider class of spacetimes generated by nonvanishing stress tensors. Then we applied this method to obtain the rotating neutral black hole, which matches the known solution found in \\cite{kerrr}. We have studied how noncommutative effects have smeared out the singularity, by calculating both the Ricci scalar and the Kretschmann invariant: they are both finite at the origin. Then we extended the analysis in \\cite{kerrr} to the thermodynamic properties of this spinning solution, providing the black hole temperature. As a second step, we have applied our method to the case of a charged and spinning object, determining the noncommutative geometry inspired Kerr-Newman black hole. We studied the geometrical properties of the solution as well as the thermodynamic ones. In both frameworks the solution shows regularity. Finally we have determined the higher-dimensional extension of these metrics, providing their Hawking temperature.  With this paper, the program of noncommutative geometry inspired black holes, initiated in 2005 with the static, neutral solution, reached its completion. As a final remark, there are some comments about future applications. On the astrophysical side, charged black holes are often regarded as remote possibilities. On the ground of Newtonian gravity consideration it has been shown that the Coulomb repulsion would prevents the formation of black holes whose charge/mass ratio exceeds $10^{-18}$. This would mean that stellar black holes could have roughly $10^2$ C, a tiny charge compared for instance to the Earth charge which is about half a million C. However due to the rotation, an intense magnetic field might occur, for the presence of superficial currents on the hole. The phenomenology of these magnetic fields would in principle deserve further investigations. On the other hand, the intrinsic quantum gravitational nature of these solutions at small length scales suggests that they could play a significant role in case of microscopic black hole phenomenology. Without entering the debate about the conjectured production of a mini black hole at the LHC, we recall that non-rotating higher-dimensional noncommutative geometry inspired black holes have been already the subject of investigations in TeV gravity phenomenology (analyses with Monte Carlo event generators can be found in \\cite{QBH}). Even in the most pessimistic scenario of the absence of quantum gravity data at the Terascale, our new solutions can play a role at least in the context of primordial black holes. Indeed it is widely accepted that the early universe could have been populated by tiny black holes, which are nowadays completely evaporated. More specifically, the conventional scenario for the evaporation of a black hole consists of four distinct phases, the spin-down phase (loss of angular momentum), the balding phase (loss of long-range fields, i.e. hair, i.e. the electromagnetic charge), the Schwarzschild phase and the Planck phase. In light of the already determined NCBH solution this general evaporation scheme has been already reviewed. For instance the  final phase is turned into a new phase, called the SCRAM \\footnote{The term SCRAM, probably the acronym for ``Safety Control Rod Axe Man'' or ``Super Critical Reactor Axe Man'' refers to an emergency shutdown of a thermonuclear reactor. The term has been extended to cover shutdowns of other complex operations or systems in an unstable state, but it also has a standard meaning ``go away quickly,'' in particular, when we address children or animals. } phase, in which the black hole cools down to a zero temperature configuration instead of experiencing a runaway increase of the Hawking temperature \\cite{NCthermo}. A further modification concerns the sequence of these phases since the discharging time depends on the number of extradimensions and on the kind of emission (brane, bulk-brane, bulk) governing both the Hawking and the Schwinger pair production mechanisms. Therefore the new NCKN solution is expected to provide the final answer about the black hole evolution from the very beginning to the very end, i.e. taking into account the yet unexplored spin-down mechanism too."
    },
    "1005/1005.0811_arXiv.txt": {
        "abstract": "We determine an empirical dense matter equation of state from a heterogeneous dataset of six neutron stars: three type I X-ray bursters with photospheric radius expansion, studied by \\\"Ozel et al., and three transient low-mass X-ray binaries. We critically assess the mass and radius determinations from the X-ray burst sources and show explicitly how systematic uncertainties, such as the photospheric radius at touchdown, affect the most probable masses and radii. We introduce a parameterized equation of state and use a Markov Chain Monte Carlo algorithm within a Bayesian framework to determine nuclear parameters such as the incompressibility and the density dependence of the bulk symmetry energy. Using this framework we show, for the first time, that these parameters, predicted solely on the basis of astrophysical observations, all lie in ranges expected from nuclear systematics and laboratory experiments. We find significant constraints on the mass-radius relation for neutron stars, and hence on the pressure-density relation of dense matter. The predicted symmetry energy and the equation of state near the saturation density are soft, resulting in relatively small neutron star radii around 11--12~km for $M=1.4\\,\\Msun$. The predicted equation of state stiffens at higher densities, however, and our preferred model for X-ray bursts suggests that the neutron star maximum mass is relatively large, 1.9--2.2$\\,\\Msun$. Our results imply that several commonly used equations of state are inconsistent with observations. ",
        "introduction": "The masses and radii of neutron stars are determined by the pressure-energy density relation (equation of state; EOS hereafter) of cold dense matter using the familiar Tolman-Oppenheimer-Volkov \\citep[TOV;][]{tolman:1939,oppenheimer39} relativistic stellar structure equations. Within tens of seconds after birth, the neutron star is cold (meaning temperature much less than Fermi energy) and deleptonized (meaning in beta equilibrium with no trapped neutrinos); as a result, for a given EOS the mass and radius of the star depend only on the central density. Inversion of the structure equations, given simultaneous mass and radius measurements, can therefore constrain the pressure-density relation, although the quality of the constraints are very sensitive to the mass and radius measurement uncertainties.  A host of observable phenomena and experimental information is becoming available~\\citep[for a recent review, see][]{Lattimer2006Neutron-Star-Ob}: on the observational side, pulsar timing, thermal emission from cooling neutron stars, surface explosions, and gravity wave emissions are some promising areas for measurements of mass and radius; on the experimental side, heavy-ion collisions, giant dipole resonances, and parity-violating electron scattering are some promising techniques for measuring the density dependence of the pressure of nuclear matter. These efforts are complementary, and determining an EOS from the many disparate measurements is a challenging task.  This paper consists of two parts. In the first, we consider simultaneous mass and radius information from astrophysical observations of X-ray bursts and thermal emissions from quiescent low-mass X-ray binaries. We critically assess the mass and radius constraints determined from X-ray bursts that may reach the Eddington limit (\\S~\\ref{s.pre-bursts})~ \\citep{vanParadijs79,vanParadijs82,Paczynski86,vanparadijs.ea:very_energetic}. Heretofore, these bursts have been interpreted assuming that the photospheric radius is equal to the stellar radius at ``touchdown'', when the effective temperature reaches a maximum~\\citep{Ozel06,Ozel09,Guver10,Guver10b,Ozel10}. If true, rather stringent constraints on mass and radius are predicted. Using the most probable values for the observed flux and angular emission area, however, does not result in real-valued masses and radii. \\citet{Ozel10} argues for rejecting much of the observed phase space on this basis and thus obtains tight constraints on masses and radii. We show that this model is not internally consistent, but that consistency can be regained by relaxing the assumption concerning the effective photospheric radius. With this modification, the inferred neutron star radii moderately increase and confidence intervals for predicted masses and radii become substantially larger than those previously quoted~\\citep{Ozel06,Ozel09,Guver10,Guver10b,Ozel10}.  We then discuss (\\S~\\ref{sec:spectral-fits}) mass and radius estimates from thermal emission from quiescent low-mass X-ray binaries (LMXBs hereafter). Neutron stars in quiescent low-mass X-ray binaries may be used to obtain angular emission areas, and thereby mass and radius information, from their quiescent thermal emission. In total, we have mass and radius constraints for six neutron stars: three bursting sources with photospheric radius expansion bursts and three quiescent neutron star transients.  While not all of the uncertainties involved in constraining the masses and radii of neutron stars are under control, it is important to {\\em quantify} the constraints on the EOS which are implied by the observations. In the second part of this paper (\\S~\\ref{sect:bayes}), we use a Bayesian analysis to combine these mass and radius constraints to determine empirical pressure-density and neutron star mass-radius relations. We include constraints to the EOS from causality, the observed minimum value of the neutron star maximum mass, and the observed maximum pulsar spin frequency. We employ a parameterized EOS (\\S~\\ref{sec:eos-parameterization}) that is compatible with laboratory measurements, and use it with a Markov Chain Monte Carlo (MC) algorithm to jointly fit these mass-radius constraints (\\S~\\ref{sec:bayes-results}). We show, for the first time, using astrophysical observations alone, that values of the nuclear parameters such as the incompressibility $K$, bulk symmetry energy $S_v$, and the density dependence of the symmetry energy $\\gamma$, all lie in ranges expected from nuclear systematics and laboratory experiments. Furthermore, our results imply that the most likely neutron star radius is relatively small, of order 11--12~km for neutron stars with masses near $1.4\\,\\Msun$, so that the predicted EOS is relatively soft in the density range 1--3 times the nuclear saturation density. Implications for the maximum mass sensitively depend on assumptions concerning X-ray burst models and are discussed in \\S~\\ref{sec:bayes-results2}. In \\S 5 we discuss astrophysical and nuclear physical implications of our results, and compare our methods and results with other studies. ",
        "conclusions": "In this paper, we have determined an empirical dense matter equation of state from a heterogeneous dataset containing PRE bursts and quiescent thermal emission from X-ray transients. Previous works \\citep{Ozel06,Ozel09,Guver10,Guver10b} have demonstrated the potential utility of PRE bursts for determining neutron star masses and radii. Their analysis assumes that the photosphere at touchdown has fully retreated to the quiescent radius, and that the high-redshift solution is favored. We argue that their model requires re-examination. First, we find no grounds on which the high-redshift solution in the case $\\rph=R$ could be favored. Second, internal consistency (i.e., the obligatory rejection of an overwhelming number of Monte Carlo realizations) implies the assumption that $\\rph=R$ is suspect and should be generalized. We then explored an alternative, that the radius of the photosphere is extended and does not recede until later in the observed burst. A larger photosphere indeed provides internal consistency without requiring strong cuts on the observed values for flux, normalization, distance, and composition.  There are other sources of systematic errors in the PRE burst model. \\citet{Boutloukos10} found that \\emph{RXTE}/PCA spectra of bursts from 4U~1820--30 were better fit by Planck or Bose-Einstein spectra rather than Comptonized spectra with large color correction factors. If the color correction factor is indeed of order unity, then the high color temperatures would indicate a locally super-Eddington flux, even in the tail of the burst \\citep{Ebisuzaki1984Are-X-ray-burst}. In addition, at near-Eddington fluxes there are other possible complications, such as a photon bubble instability \\citep{Hsu97} that may impact the observed properties of a PRE X-ray burst.  A recent analysis of PRE bursts in 4U~1724--307 by~\\citet{Suleimanov10} also found that the photosphere is somewhat extended at touchdown, $\\rph \\approx 2 R$. Using a new model of the atmosphere (the details of which are not yet published), they obtain masses and radii that are both larger than we would have found. This result is principally due to a larger value of the color correction factor, $f_c$, predicted by their model atmosphere.  It is 10--15\\% larger than what one would obtain from the model of \\citet{Madej04} for the same values of $F/\\Fedd$, surface gravity $g$, and composition.  In this case, we find values of $\\alpha\\propto f_c^{-2}$ which are about 25\\% larger and $R_\\infty\\propto f_c^2$ that are about 25\\% smaller than \\citet{Suleimanov10} did.  Assuming $\\rph\\gg R$, the predicted mass scales with $f_c^2$.  When we analyse 4U~1724--307 with the lower value of $f_c$, we get masses and radii that are consistent with those from the other three PRE burst sources. This highlights an important avenue for future work, namely that a clearer understanding of the atmosphere and, in particular, $f_c$ are essential.  Further progress in using PRE bursts to constrain neutron star masses and radii clearly requires better models of the spectral evolution during X-ray bursts. It is important to note that increases in the precision of mass and radius estimates are not absolutely necessary, as we have shown that existing errors, large as they are, do not inhibit placing interesting limits to the equation of state when combined with other observations. It is very important, however, to resolve the uncertainty regarding the location of the photospheric radius at ``touchdown'' and to characterize systematic uncertainties, including potential correlations between $\\Ftd$, $A$, $f_c$, and $X$ in the spectral models. Also, should a larger range in $f_c$ be required, then the uncertainties on masses and radii will increase accordingly.  Our results imply that the EOS in the vicinity of the nuclear saturation density is relatively soft. As shown by \\citet{Lattimer01}, this is primarily due to the weak dependence of the nuclear symmetry energy with density. This conclusion is robust with respect to variations in how PRE burst sources are modeled, and is strengthened by the small values of $R_\\infty$ deduced for the globular cluster LMXB sources in M13 and $\\omega$ Cen. This conclusion has immediate and important ramifications for laboratory nuclear physics, in particular for the scheduled parity-violating electron scattering measurement of the neutron skin thickness of $^{208}$Pb~\\citep{Michaels00,Horowitz01b}. In the context of the liquid droplet model, the ratio of the surface and volume symmetry coefficients can be expressed as~\\citep{Steiner05} \\begin{equation} {S_s\\over S_v}={9E_{s0}\\over K}{I_2\\over I_1}, \\label{eq:sssv} \\end{equation} where $E_{s0}\\simeq19$ MeV is the surface energy coefficient for symmetric nuclei.  The integrals $I_1$ and $I_2$ for the schematic EOS described by equation~(\\ref{eq:ldeos}) are given by \\begin{eqnarray} I_1 &=& \\int_0^1\\sqrt{u}(1-u)\\sqrt{1-a(1-u)}~du,\\cr I_2 &=& \\int_0^1\\sqrt{u}\\left[{S_v\\over S_ku^{2/3}+(S_v-S_k)u^\\gamma}-1\\right] {1\\over1-u}{1\\over\\sqrt{1-a(1-u)}}~du\\,, \\label{eq:i1i2} \\end{eqnarray} where $a=K^\\prime/(9K)$. The predicted neutron excess in the center of a nucleus $(N,Z)$ is then \\begin{equation} \\delta={N-Z\\over N+Z}\\left(1+{S_s\\over S_vA^{1/3}}\\right)^{-1}, \\label{eq:delta} \\end{equation} and the neutron skin thickness is \\begin{equation} \\delta R=\\sqrt{3\\over5}{2r_0\\over3}{S_s\\over S_v}{\\delta\\over1-\\delta^2}\\,, \\label{eq:dr} \\end{equation} where $r_0=(4\\pi n_0/3)^{-1/3}$. With a value of $\\gamma=0.26$ and $a=-0.141$ and errors as established in Table~4, we deduce $S_s/S_v\\simeq1.5\\pm0.3$ and $\\delta R(^{208}\\mathrm{Pb})\\simeq0.15\\pm0.02$ fm, a value at the lower end of expectations.  We also conclude that in our preferred model, in which the photospheric radius is not restricted to be the stellar radius, the EOS at high densities is stiff enough to support a maximum mass of order $2\\, \\Msun$ or greater. This result is supported by the simultaneous observations of LMXB sources with both small and large values of $R_\\infty$, if one rejects the possibility of the existence of neutron stars with masses less than about $1\\,\\Msun$. The latter condition seems to be borne out by models of massive star supernovae, which strongly suggest that proto-neutron stars are born with high trapped lepton fractions and moderate entropies. Hydrostatic stability requires, in this case, that protoneutron stars are bound only if their mass exceeds about $1\\,\\Msun$~\\citep{Strobel99}, ruling out the existence of cold, catalyzed neutron stars of smaller masses. Had this assumption been used as a prior condition in our analysis, the LMXB sources with small values of $R_\\infty$ would further support solutions with EOS parameters favoring large maximum masses.  Obviously, more neutron star observations with mass and radius constraints would enable one to improve our constraints. It is a particular advantage of our methodology that new observations can be integrated into the formalism easily, if estimates of a two-dimensional probability distribution of mass and radius can be made. As pointed out by~\\citet{Heinke06}, it is particularly important that atmosphere models treat the surface gravity self-consistently. In particular, spectral features would be very constraining, and although the most recent result from~\\citet{Cottam02} has not been verified in longer observations \\citep{Cottam2008The-Burst-Spect}, future determinations of the surface gravity may provide strong constraints. \\citet{Ho09} also demonstrated that constraints on the mass and radius of the Cas A supernova remnant may be obtained. These constraints are roughly consistent with our $\\rph \\gg R$ results described above.  Although we did not include the estimate of the mass and radius ($M=1.7\\pm0.3\\, \\Msun$ and $R=11.5\\pm1.2\\,\\mathrm{km}$) of RX~J1856--3754 \\citep{Pons02,Walter10} in our baseline fit, we found them to be remarkably consistent with those of the other 6 neutron stars. All of the results we obtained for the parameters of the EOS, the pressure versus energy density curve, the mass versus radius curve, our estimate of the maximum mass, and the predicted individual masses and radii for all 6 other sources were unchanged to within 1-$\\sigma$. For example, the most probable radius of a $1.4\\,M_\\odot$ star would change by less than $0.01\\,\\mathrm{km}$ and the most probable pressure at an energy density of $1000\\,\\mathrm{MeV\\,fm^{-3}}$ would increase by 0.03\\%.  In addition to X-ray bursts, quiescent low-mass X-ray binaries, and thermally-emitting isolated neutron stars, other potential methods for determining neutron star masses and radii exist. These include neutron star seismology~\\citep{Samuelsson07,Steiner09}, pulse profiles in X-ray pulsars~\\citep{Leahy2009Constraints-on-,Morsink2009Multi-epoch-Ana}, and moment of inertia measurements in relativistic binaries~\\citep{Lattimer05}. Gravitational wave signals of neutron star mergers may also provide significant constraints~\\citep{Ferrari10}. Measurements of the thickness of the neutron star crust, which controls the crust cooling of transiently accreting LMXBs~\\citep{Shternin2007Neutron-star-co,Brown2009Mapping-Crustal} as well as the distribution of observed cooling curves compared to a minimal cooling model \\citep[e.g.,][]{Page09}, could also constrain neutron star masses and radii.  The Bayesian analysis in Section~\\ref{sect:bayes} is a novel procedure for combining mass and radius constraints from disparate objects to form new constraints on the mass versus radius curve and the EOS of dense matter. The construction of the EOS from astrophysical observations has also been recently addressed in~\\citet{Read09},~\\citet{Ozel09b}, and \\citet{Ozel10}. \\citet{Read09} showed that piecewise polytropes with relatively few parameters could effectively describe more sophisticated models for a high-density EOS, but confined attention to constraints stemming from observations limiting the neutron star maximum mass and maximum spin rate. They did not attempt reconstruction of the EOS from simultaneous mass and radius measurements. We therefore obtain stronger constraints on the EOS, but have utilized observational data which contains significant systematic uncertainties. \\citet{Ozel09b} examined the constraints on the EOS obtained from a synthetic data set obtained with simultaneous mass and radius measurements of either 5\\% or 10\\% accuracy for three separate objects; they found that although individual EOS parameters were difficult to estimate precisely, significant correlations between them could be obtained. The Jacobian technique employed in \\citet{Ozel09b} is a simple approximation of our full Markov Chain Monte Carlo method. That Jacobian technique requires that the number of EOS parameters and the number of neutron stars are equal, and thus provides only an incomplete marginalization over the EOS parameters.  One could imagine various alternatives to the statistical procedure in Section~\\ref{sect:bayes}. One possibility is the use of prior distributions to represent either astrophysical or nuclear input. Neutron star masses in some double-pulsar systems are well-measured, and could provide a prior mass distribution. It is not clear, however, that either low-mass X-ray binaries or isolated neutron stars follow this same initial mass function. As mentioned above, there are strong theoretical reasons to suspect that neutron stars cannot be formed with less than about $1.0\\,\\Msun$. In the analysis above, we chose $0.8\\,\\Msun$ to ensure that the boundary does not interfere with results near the proto-neutron star minimum mass. There are several nuclear physics observables which can be used to constrain the EOS through measurements of properties like the compressibility and the symmetry energy. Several authors have also computed the EOS of neutron matter directly from nucleon-nucleon interactions~\\citep{Tolos07,Hebeler09}, and we have found that our predicted EOS is consistent with these studies. In the context of this work, information from calculations of the neutron matter EOS, neutron skin thicknesses, the surface symmetry-volume symmetry energy correlation observed from mass formula fits~\\citep{Steiner05}, information from giant dipole resonances, and so forth, could also provide constraints for the schematic EOS parameters described above. We have chosen not to use this information in this work, in part to ensure that our results are free from extra model dependencies. Future work on implementing more nuclear physics input into the analysis of the observational data is certainly warranted.  One might also question, for the high-density EOS, whether or not our prior distributions for the polytropic indicies ought to be treated as uniform, as a uniform distribution in the polytropic index is different than, for example, a uniform distribution in the pressure at any fixed energy density, or in the polytropic exponent. Finally, one could consider reformulating the model directly in terms of the observables like flux and distance instead of the ``two-step'' procedure we have used above which uses a MC simulation to generate masses and radii from the observables for each neutron star, and then afterwards performs a Bayesian analysis from the output of these initial MC results. This alternative would strongly disfavor $\\rph=R$, since so many MC realizations must be rejected in this picture.  It is important to note that the constraints we have obtained are a guide, but will likely require revision in the future. Because we only used 6 mass and radius measurements a single additional measurement could have a significant impact on our results. Alternatively, if one of the 6 input probability distributions used above changes significantly, because, for example, of a new understanding of systematic uncertainties, our final constraints on the EOS would change accordingly. We have already noted that there is significant tension on our results from assumptions about the photospheric radius of X-ray bursts, and also from the fact that the neutron star in M13 has a small value of $R_{\\infty}$ while X7 in Terzan 5 has a much larger value of $R_{\\infty}$.  It is possible that there exist extreme models which live in very small regions of parameter space and are not fully sampled in this work. One example of such an EOS would be those which exhibit the ``twinning'' phenomenon, i.e. the presence of a turning point in the mass vs. radius curve which admits a new family of compact neutron stars~\\citep{Glendenning00,Schaffner-Bielich02} Such solutions should be addressed in future work.  \\citet{Ozel10} claim that constraints from the PRE burst sources imply that equations of state which contain only nucleonic degrees of freedom are inconsistent with data. We disagree with their conclusion. Although we find the EOS must be quite soft at moderate densities, we do find many nucleonic models which are consistent with the data. More importantly, however, we find that the conclusion of extreme softening evaporates if we use a slightly different model for the PRE burst sources that accounts for systematic uncertainties. The inclusion of mass and radius data from other neutron star sources supports our interpretation of the PRE burst sources. While our results do not rule out a phase transition at supernuclear densities, extreme softening of the EOS is not compatible with observations. Rather, the implication is that the maximum mass is likely large, greater than $1.8\\,\\Msun$."
    },
    "1005/1005.2025_arXiv.txt": {
        "abstract": "A cooperative observation with Hida observatory and Hinode satellite was performed on an emerging flux region. The successive Ca\\emissiontype{II} K spectro-heliograms of the emerging flux region were taken by the Domeless Solar Telescope of Hida observatory. Hinode observed the emerging flux region with Ca\\emissiontype{II} H and Fe\\emissiontype{I} Stokes IQUV filtergrams. In this study, detailed dynamics and temporal evolution of the magnetic flux emergence was studied observationally. The event was first detected in the photospheric magnetic field signals. 3 minutes later, the horizontal expansion of the dark area was detected. And then, 7 minutes later than the horizontal expansion, the emerging loops were detected with the maximal rise speed of 2.1 km/s at chromospheric heights. The observed dynamics of emerging magnetic flux from the photosphere to the upper chromosphere is well consistent with the results of previous simulation works. The gradual rising phase of flux tubes with a weak magnetic strength was confirmed by our observation. ",
        "introduction": "An emerging flux region (EFR) is the early stage of an active region where magnetic flux loops emerge from underneath the solar photosphere \\citep{Bru69, Zir72}. With H$\\alpha$ line center, typical EFRs show a pair of bright regions and dark loops between them. The dark loop and the cluster of them are called as the arch filament and the arch filament system (AFS), respectively \\citep{Bru67}. The typical size of arch filaments is about 30,000 km \\citep{Bru67}. The lifetime of an arch filament is 10-30 minutes and the rise velocity is from $-10$ to $-15$ km s$^{-1}$ \\citep{Bru67, Cho88}. The both ends of an arch filament are called as footpoints. The footpoints appears as plage regions in H$\\alpha$ line center image. At the footpoints of an arch filament there are strong downflows with 35-50 km s$^{-1}$.  Although Ca\\emissiontype{II} H and K lines have been used for solar observation, there are few studies about the Ca\\emissiontype{II} H and K spectroscopic observation of EFRs. \\citet{Bru69} observed an AFS with the Ca\\emissiontype{II} K line filtergram and reported that the arch filaments were less pronounced in the Ca\\emissiontype{II} K than in H$\\alpha$. Since Ca\\emissiontype{II} H and K lines show double reversed spectral profiles, it is difficult to determine the Doppler shift of their components. \\citet{Zwa85} performed a spectroscopic observation on an EFR and measured the Doppler shift of Ca\\emissiontype{II} H3 absorption core, which showed about 7 km s$^{-1}$ rising motion at the apex of arch filaments and at least 40 km s$^{-1}$ downflows at the footpoints. They also measured the Dopplershift of the Fe\\emissiontype{I} lines at 6302.5 {\\AA} and 5691.5 {\\AA} (both are photospheric lines) to be 1.5 km s$^{-1}$ downflows at the footpoints. The spectro-heliogram of a fairly large EFR with high spatial resolution were presented by \\citet{Bal01}. He picked up the spectra of the footpoints of arch filaments and found that there were additional absorptions on the K$_2$ red peak (K$_{2R}$). The Doppler shift of the additional absorption was of the order 40-50 km s$^{-1}$. He interpreted them as due to the formation of the shocks at the footpoints.  There is a threshold flux for an AFS. \\citet{Har73} estimated a minimum total flux of $1.5 \\times 10^{20}$ Mx to form an AFS. \\citet{Cho87} found the threshold flux to be $0.5 \\times 10^{20}$-$1.0 \\times 10^{20}$ Mx for an AFS formation.  As the observation technologies advanced, smaller scale emergences of magnetic flux have been observed. Especially the Solar Optical Telescope (SOT: \\cite{Ich04, Tsu07, Sue07, Shi07}) on-board Hinode Satellite has revealed the microscopic dynamics and faint structures of small magnetic phenomena. \\citet{Cen07} observed a granular size ($\\sim$\\timeform{2''}) EFR in the quiet-Sun internetwork using the spectro-polarimeter of SOT. \\citet{Ots07} investigated the temporal evolution of a small-scale EFR with the Ca\\emissiontype{II} H filtergram, the Fe\\emissiontype{I} Stokes I and V filtergram. He found the flux tube expanding laterally in the photosphere and the lower chromosphere with a speed of 3.8 km s$^{-1}$. A farther study on small-scale EFR was performed by \\citet{Gug08}, who showed the relation between the Ca\\emissiontype{II} H intensity and the magnetic flux of the footpoints. He measured the total flux of the EFR to be the order of $10^{18}$ Mx.  On the other hand, many numerical simulation of EFRs have been performed \\citep{Shi89, Mat92, Fan01}. \\citet{Mat93} simulated the three-dimensional magnetohydrodynamics (MHD) of the emerging magnetic flux. In his study, magnetic loops with less than $10^{19}$ Mx can not expand into the corona. \\citet{Mag01}, \\citet{Noz05} and \\citet{Mur06} also performed the MHD simulations of EFRs with less magnetic flux, which showed lateral expansions and less uprising motions.  In this paper we report the spectroscopic investigation of the small-scale EFR obtained by Hida-Hinode cooperative observation (HOP12). Firstly, from a morphological standpoint, we compare the images of the EFR observed in Hida and by Hinode. Next we perform a double Gaussian fitting on the Ca\\emissiontype{II} K line spectra to obtain the Doppler shift of K$_{2,3}$ components. The temporal evolutions of the obtained physical parameters of the EFR are compared each other, then we figure out the dynamics of the emergence event. At the end of this paper we discuss about the results of our observational study comparing with previous MHD simulations. ",
        "conclusions": "We observed an EFR with Hinode/SOT and DST/VS. From the successive images and the temporal variations of the physical parameters observed by SOT and DST, we can summarize that the scenario for the emergence of the EFR as follows: \\begin{enumerate} \\item The small preceding Ca filament appeared in SOT filtergram at 02:50 UT without any Stokes I darkening and horizontal field increase. The width of the preceding Ca filament reached the maximum of \\timeform{0.62''} at 02:53 UT with the lateral expansion speed of 2.7 km s$^{-1}$. \\item The Doppler shift of Ca\\emissiontype{II} K$_2$ showed the blue shift from 02:56 UT with the acceleration of $-8 \\times 10^{-3}$ km s$^{-2}$. The observed maximum velocity for the preceding filament was $-2.2$ km s$^{-1}$ at 03:00 UT. \\item The horizontal field increased at 02:56 UT and lasted until 03:18 UT associated with the main Ca filament. Just simultaneously, the dark granular lane in the Stokes I image started to appear. \\item The main Ca filament appeared in SOT filtergram at 02:59 UT, 3 minutes after the horizontal field emergence. The filament grew with the elongation speed of $\\sim$ 4 km s$^{-1}$ and the expansion speed of 2.9 km s$^{-1}$. The main filament reached its maximum width of \\timeform{1.03''} at 03:08 UT. The K$_2$ Doppler velocity was about $-1.0$ km s$^{-1}$ for this 4 minutes period. Ca\\emissiontype{II} K$_3$ image did not show a prominent darkening. \\item Then the Doppler shift in Ca\\emissiontype{II} K$_2$ corresponding to the main filament exhibited the acceleration of $-3 \\times 10^{-3}$ km s$^{-2}$ from 03:10 UT. The maximum velocity of $-2.1$ km s$^{-1}$ were observed at 03:15 UT for the main filament. \\end{enumerate}  We consider that the time-lag of 10 minutes between the appearance of the horizontal field (02:56 UT) and the K$_{2V}$ darkening (03:06 UT) is due to the formation height difference. Assuming that the K$_2$ formation height to be about 700 km \\citep{Ver81}, the average rise velocity is calculated as about $-1.2$ km s$^{-1}$. The estimated rise velocity is similar to the observed K$_2$ Doppler velocities, which varied between $-1$ km s$^{-1}$ and $-2$ km s$^{-1}$. These observed values are consistent with the result of numerical simulations very well.  The temporal difference between the K$_{2V}$ darkening and the beginning of the acceleration is interpreted as below: When arch filament rises above the photosphere, the flux tube expands both laterally and vertically. However, if the total flux is not enough, the flux tube can not rise but expands laterally \\citep{Mat93, Noz05}. Then the laterally expanded flux tube gradually rises \\citep{Mag01, Mur06} and will wait for the buoyancy instability to develop. The velocity of the gradually rising phase is in the range of $-0.4$ km s$^{-1}$ to $-1.7$ km s$^{-1}$ \\citep{Mur06}. This value is consistent with the constant velocity of about $-1$ km s$^{-1}$, which was observed in this study. We estimated the total flux of the arch filament using SOHO/MDI data. The magnetic field strength is less than 100 G at the footpoint of the arch filament. The size of footpoint is about \\timeform{2''}, so the upper limit of the total flux is about $2\\times10^{18}$ Mx. This value is smaller than the threshold mentioned by \\citet{Mat93}. So this emerging flux is thought to be unable to reach the corona.  The mean upward velocity ($-1$ km s$^{-1}$) of the arch filament measured in this study is fairly smaller than the Doppler velocity of K$_3$ component obtained by \\citet{Zwa85}. With this velocity, during the rising phase of 20 minutes the gas can not reach the formation height of K$_3$ ($\\sim 2000$ km). Actually, our K$_3$ spectro-heliogram image did not show any prominent darkening. So we suggest that the flux tube with less magnetic field strength can gain too small buoyancy to rise to the upper chromosphere  and the corona, as was indicated by the simulation of \\citet{Mur06}.  Finally we give a comment on the time difference of 3 minutes between the darkenings observed in Stokes I and in Ca\\emissiontype{II} H of Hinode. This value is much smaller than the time-lag between the appearances of the Stokes I darkening and the K$_{2V}$ darkening (10 minutes). Since the photospheric dark features may be detected in Hinode Ca\\emissiontype{II} H images due to the broadness of the filter pass band, the rising loops can be seen in Ca\\emissiontype{II} H BFI images earlier than the arrival to chromospheric heights. In our study, we could detect the real rising motion of loops with the help of spectroscopic analysis.  In conclusion, we have first observationally detected the detailed evolution of the magnetic flux emergence of less magnetic flux than the active EFRs and confirmed that the observed evolution is consistent with those simulated and predicted by the previous theoretical works. \\\\  We express our gratitude to the referees whose comments helped us to revise the paper. We are grateful to all the staff members of the Kwasan and Hida Observatories, Kyoto University for their guidance of the Domeless Solar Telescope and for the fruitful discussions with them. This work was supported by the Grant-in-Aid for the Global COE Program ``the Next Generation of Physics, Spun from Universality and Emergence'' from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan, by the Grant-in-Aid for Creative Scientific Research The Basic Study of Space Weather Prediction (17GS0208, PI: K. Shibata) from the Ministry of Education, Science, Sports, Technology, and Culture of Japan, and by the grant-in-aid from the Japanese Ministry of Education, Culture, Sports, Science and Technology (19540474). Hinode is a Japanese mission developed and launched by ISAS/JAXA, collaborating with NAOJ as a domestic partner, NASA and STFC (UK) as international partners. Scientific operation of the Hinode mission is conducted by the Hinode science team organized at ISAS/JAXA. This team mainly consists of scientists from institutes in the partner countries. Support for the post-launch operation is provided by JAXA and NAOJ (Japan), STFC (U.K.), NASA, ESA, and NSC (Norway)."
    },
    "1005/1005.3812_arXiv.txt": {
        "abstract": "We use a sample of newly-discovered globular clusters from the Pan-Andromeda Archaeological Survey (PAndAS) in combination with previously-catalogued objects to map the spatial distribution of globular clusters in the M31 halo. At projected radii beyond $\\approx30$ kpc, where large coherent stellar streams are readily distinguished in the field, there is a striking correlation between these features and the positions of the globular clusters. Adopting a simple Monte Carlo approach, we test the significance of this association by computing the probability that it could be due to the chance alignment of globular clusters smoothly distributed in the M31 halo. We find the likelihood of this possibility is low, below $1\\%$, and conclude that the observed spatial coincidence between globular clusters and multiple tidal debris streams in the outer halo of M31 reflects a genuine physical association. Our results imply that the majority of the remote globular cluster system of M31 has been assembled as a consequence of the accretion of cluster-bearing satellite galaxies. This constitutes the most direct evidence to date that the outer halo globular cluster populations in some galaxies are largely accreted. ",
        "introduction": "It has long been suspected that a significant fraction of the Milky Way's globular clusters formed in smaller `proto-galactic fragments' that were subsequently accreted into the Galactic potential well. First proposed in the seminal paper by \\citet{searle:78}, there has since been gradual accumulation of indirect evidence in support of this hypothesis -- modern data suggest that the abundances, ages, velocities, horizontal-branch morphologies and sizes of many globular clusters at Galactocentric radii $\\ga10$ kpc are consistent with an external origin \\citep[e.g.,][]{zinn:93,mackey:04,marinfranch:09}. Even so, it has proven problematic to unambiguously identify indvidual clusters as having been accreted into the Galaxy. The only {\\it direct} observation of this process is the disrupting Sagittarius dwarf \\citep{ibata:94} which is depositing at least five globular clusters into the Milky Way halo \\citep{dacosta:95,martinezdelgado:02,bellazzini:03}; more controversially, the putative Canis Major dwarf may also be responsible for several new arrivals \\citep{martin:04}. The extent to which the observed properties of sub-groups within the Milky Way globular cluster system reflect the assembly history of the Galactic halo thus remains a critical unresolved question.  As the nearest large spiral galaxy, M31 is an attractive alternative target for studying this problem. Its globular cluster system is several times larger than that of the Milky Way, and observational investigation of its halo is less prone to the vagaries of projection and extinction that plague Galactic surveys. It is known that globular clusters projected near its central regions exhibit some evidence for sub-clustering in position-velocity space that may signal an accretion origin \\citep{ashman:93,perrett:03}, although interpretation is difficult because of the complex nature of the inner M31 system. Potentially less confusing are halo regions at projected radii $R_{{\\rm p}}\\ga15$ kpc, where dynamical times are also longer; however, it is only relatively recently that these remote areas have been targeted by deep wide-field surveys. Various such studies have shown the M31 halo to be littered with coherent tidal debris features indicative of one or more accretion events \\citep{ferguson:02,ibata:07}, and have also facilitated the discovery of significant samples of remote M31 globular clusters \\citep{huxor:08} so that it is now possible to begin assessing how these objects relate to the stellar halo in this galaxy.  \\citet{huxor:10} have derived the first radial surface-density profile for M31 globular clusters to extend into the far outer halo ($R_{{\\rm p}}\\approx100$ kpc). Their profile exhibits a distinct flattening beyond $R_{{\\rm p}}\\approx30$ kpc, very similar to that observed for the metal-poor field halo, and has been interpreted as evidence that accretion processes have played a role in building up both components.  In this Letter we use new results from the {\\it Pan-Andromeda Archaeological Survey} \\citep[PAndAS;][]{mcconnachie:09} to provide the most extensive map to date of globular clusters in the M31 halo, and explore the implications for the origin of this system. PAndAS is an ongoing large program on the Canada-France-Hawaii Telescope, utilising the MegaCam imager to obtain a deep panoramic view of M31 and M33. First-semester imaging and data reduction was completed in mid-2009, revealing in exquisite new detail the abundance of low surface brightness substructure present in the M31 halo \\citep[][see also Fig. 1]{mcconnachie:09}. ",
        "conclusions": "Put together, our results provide strong evidence that globular clusters in the outer halo of M31 are anisotropically spatially distributed, and preferentially associated with underlying tidal debris features. Of the $61$ clusters in the PAndAS footprint with $R_{{\\rm p}}\\ga30$ kpc, at least $27$ lie on the major substructures outlined in Section \\ref{ss:streams}; if the complex major-axis region is also included, this number rises to $\\approx37$. A further $13$ objects are members of the two cluster overdensities described in Section \\ref{ss:clustergroups}, leaving just $11$ of the sample unaccounted for. Figure 1 reveals only a handful of these lie away from field overdensities altogether.  These numbers imply that the majority ($\\ga80\\%$) of the outer globular cluster system of M31 has been built up via the accretion of satellite host galaxies. This fraction matches closely that inferred for the outer Galactic system \\citep[e.g.,][]{mackey:04,forbes:10}. Our work provides a striking direct illustration of the \\citet{searle:78} paradigm; along with the disrupting Sagittarius dwarf, it represents the most clear-cut observation to date of the assembly of a globular cluster system in action, and does so on a grand scale -- across an entire galactic system and a huge population of clusters.  We note that all but one of the extended clusters in our sample with $R_{{\\rm p}}\\ga30$ kpc are projected onto stellar substructure or are members of a cluster overdensity. This observation is consistent with the idea that these puzzling objects may predominantly originate in lower-mass galaxies \\citep[e.g.,][]{dacosta:09}.  Intriguingly, streams without clusters are seemingly in the minority in the outer M31 halo. With the caveat that there may be very low-luminosity features falling below the PAndAS surface-brightness limit, this could imply that most substructure in the outer M31 halo is due to the accretion of just a few larger cluster-bearing satellites. This would be in qualitative agreement with the results of cosmologically-motivated simulations, which suggest that the haloes of large spiral galaxies are predominantly assembled via the accretion of a handful ($\\la5$) of significant progenitors \\citep[e.g.,][]{bullock:05,cooper:10}.  Owing to their recent discovery, little is yet known about most of the remote M31 clusters described in this paper. However, a few have published colour-magnitude diagrams and/or spectra, which reveal them as almost exclusively metal-poor with $-2.3\\la[$Fe$/$H$]\\la-1.6$ \\citep{mackey:06,mackey:07,mackey:10, alvesbrito:09}. Others exhibit a small dispersion in integrated colour, further supporting this assertion \\citep{huxor:10}. This might imply that a significant fraction of metal-poor globular clusters in large galaxies are accreted objects, as suggested by a variety of models \\citep[e.g.,][]{cote:98,cote:00,prieto:08,muratov:10}. Our present results raise the exciting imminent prospect of characterizing, for the first time, the individual properties of multiple accreted families of globular clusters in a galactic halo."
    },
    "1005/1005.3215_arXiv.txt": {
        "abstract": "{The existence of dusty debris disks around a large fraction of solar type main-sequence stars, inferred from excess far-IR and submillimetre emission compared to that expected from stellar photospheres, suggests that leftover planetesimal belts analogous to the asteroid- and comet reservoirs of the Solar System are common.} {In order to detect and characterise cold extended dust originating from collisions of small bodies in disks, belts, or rings at Kuiper-Belt distances (30--50\\,AU or beyond) sensitive submillimetre observations are essential. Measurements of the flux densities at these wavelengths will extend existing IR photometry and permit more detailed modelling of the Rayleigh-Jeans tail of the disks spectral energy distribution (SED), effectively constraining dust properties and disk extensions. By observing stars spanning from a few up to several hundred Myr, the evolution of debris disks during crucial phases of planet formation can be studied.} {We have performed 870\\,$\\mu$m observations of 22 exo-Kuiper-Belt candidates, as part of a Large Programme with the LABOCA bolometer at the APEX telescope. Dust masses (or upper limits) were calculated from integrated 870\\,$\\mu$m fluxes, and fits to the SED of detected sources revealed the fractional dust luminosities $f_{\\mathrm{dust}}$, dust temperatures $T_{\\mathrm{dust}}$, and power-law exponents $\\beta$ of the opacity law.} {A total of 10 detections with at least 3$\\sigma$ significance were made, out of which five (HD\\,95086, HD\\,131835, HD\\,161868, HD\\,170773, and HD\\,207129) have previously never been detected at submillimetre wavelengths. Three additional sources are marginally detected with $>2.5\\sigma$ significance. The best-fit $\\beta$ parameters all lie between 0.1 and 0.8, in agreement with previous results indicating the presence of grains that are significantly larger than those in the ISM. From our relatively small sample we estimate $f_{\\mathrm{dust}}$\\,$\\propto$\\,$t^{-\\alpha}$, with $\\alpha$\\,$\\sim$\\,0.8--2.0, and identify an evolution of the characteristic radial dust distance $R_{\\mathrm{dust}}$ that is consistent with the $t^{1/3}$ increase predicted from models of self-stirred collisions in debris disks.} {} ",
        "introduction": "The existence of warm dust around solar type main-sequence stars is implied from infrared (IR) excess emission (above the stellar photospheric contribution) discovered by the \\emph{Infrared Astronomical Satellite} \\citep[IRAS, e.g.][]{Aumann1984,Neugebauer1984}, and studied by the \\emph{Infrared Space Observatory} \\citep[ISO, e.g.][]{Kessler1996}, and \\emph{Spitzer} \\citep[e.g.][]{Werner2004}. Since such dust has a limited lifetime its existence suggests the presence of larger asteroidal and/or cometary bodies which through collisions continuously replenish the dusty debris disk. Excess emission at wavelengths of 25\\,{$\\mu$}m and shorter comes predominantly from hot dust ($>$300\\,K), located within 1--5\\,AU of the star (similar to the asteroidal belt in our Solar system), while an excess at 25--100\\,{$\\mu$}m probes cooler dust ($\\sim$100--300\\,K) located at 5--10\\,AU. However, in order to detect cold and very extended dust, originating in Kuiper-Belt analogs (on the order of 100\\,AU from the star), observations at submillimetre (submm) wavelengths are needed. Although very few older stars show evidence of hot dust, likely due to clearing by planets \\citep[e.g.][]{Beichman2006a,Meyer2008}, cool debris disks appear more common and many should be detectable at submm wavelengths based on extrapolation from observed far-IR fluxes.  The potential correlation of cold dust and planetesimal belts with the formation of exo-planets in these systems is another incentive for larger submm surveys of nearby solar-type main-sequence stars. \\citet{Beichman2006b} and \\citet{Bryden2006} reported \\emph{Spitzer} MIPS observations of a large sample of nearby F, G, and K main-sequence stars, and detected an overall excess of 70\\,{$\\mu$}m emission towards 13\\% of the stars. From a larger sample of solar-type stars, spanning ages between 3\\,Myr and 3\\,Gyr, studied in the \\emph{Spitzer} Legacy programme Formation and Evolution of Planetary Systems (FEPS), showed a total 70\\,{$\\mu$}m excess rate of 7\\% with high (50 times the photosphere) fractional excesses at ages between 30 and 200\\,Myr, and significantly lower fractional excess for older stars \\citep{Carpenter2009,Hillenbrand2008a}. There was no clear trend for the temporal evolution of 70\\,{$\\mu$}m excess in FEPS, however \\citet{Su2006} estimate a decay time of $\\sim$400\\,Myr for A type stars. At submm and mm wavelengths 3\\% of the FEPS stars show evidence of continuum emission, with a suggested decrease in dust masses and/or changes in the grain properties for stellar ages of 10--30\\,Myr \\citep{Carpenter2005}.  An additional advantage in making observations in the submm region is that measured integrated fluxes are directly proportional to the temperature and mass of the disk, due to the fact that they (in most cases) sample the Rayleigh-Jeans tail of the spectral energy distribution (SED), and that the disk can be assumed to be optically thin at these wavelengths. At the same time, the thermal radiation in the submm is dominated by large and cool dust grains, thus providing a better estimate of the total dust mass (although still giving just a lower limit). By combining submm observations with IR-photometry a fit to the SED can be made, yielding constraints on the temperature and radial extent of spatially unresolved disks.  In this paper we present the first results of a Large Programme submm-survey with the \\emph{Large APEX Bolometer CAmera} (LABOCA) at the \\emph{Atacama Pathfinder EXperiment} (APEX) telescope, targeting stars of spectral type B to M, belonging to nearby stellar associations with ages ranging between 10 and 100\\,Myr. While the main aim of the study is to investigate disk evolution by relating disk parameters and grain properties to fundamental stellar parameters (metallicity, mass, and age), we also hope to spatially resolve extended dust disks around some of the most nearby stars. The observational scheme is designed to reach $\\sim$2\\,mJy/beam root-mean-square (RMS) sensitivity, enabling the detection of debris disk masses as low as the mass of the Moon \\citep[comparable to the Solar System's Kuiper-Belt mass of 3\\,$M_{\\mathrm{Moon}}$;][and references therein]{Iorio2007}, and to make observations of exo-Kuiper-Belts out to distances of about 50--100\\,pc. A successful precursor study of the $\\sim$12\\,Myr-old $\\beta$ Pictoris Moving Group clearly detected 2 out of 7 stars \\citep{Nilsson2009}. ",
        "conclusions": "The most important results from our 870\\,$\\mu$m observations of 22 exo-Kuiper-Belt candidates can be summarised as follows: \\begin{itemize} \\item[\\textbullet] Out of the observed sample of far-IR excess stars we detected $\\sim$45\\% (10 out of 22) with at least a 3$\\sigma$ significance. Five of the detected debris disks (HD\\,95086, HD\\,131835, HD\\,161868, HD\\,170773, and HD\\,207129) have previously never been seen at submm wavelengths. These findings increase the currently known exo-Kuiper-Belts (inferred from cold extended dust disks detected in the submm) from 29 to 34. The sensitivity of current submm observations with LABOCA at APEX and SCUBA at JCMT are reaching the background extragalactic confusion limit, as evident by the abundance of significant off-center flux density peaks in acquired submm maps. Future ALMA observations will be essential to discriminate between resolved disks and background submm galaxies. \\item[\\textbullet] We perform SED modelling of available photometric data using a simple modified blackbody to make a $\\chi^{2}$ minimization fit to the disk excess emission in order to determine the dust temperature $T_{\\mathrm{dust}}$ and power-law exponent $\\beta$ of the opacity law. Resulting $\\beta$ parameters are all between 0.1 and 0.8, suggesting grains significantly larger than those in the ISM, in agreement with previous results for debris disks \\citep[e.g.][]{Najita2005,Nilsson2009,Roccatagliata2009}. However, it should be noted that the presence of colder secondary dust components in addition to the fitted disk could be responsible for some abnormally low $\\beta$ values. The measured 870\\,$\\mu$m flux of HD\\,95086 undeniably requires a second disk component. \\item[\\textbullet] Besides finding the dust mass of detected sources (together with upper limits on undetected ones) derived from the integrated 870\\,$\\mu$m flux density, we use the best-fit SED parameters to calculate the fractional dust luminosity $f_{\\mathrm{dust}}$ and the characteristic radial dust distance $R_{\\mathrm{dust}}$ (for 1$\\mu$m sized grains) of the disks. We present plots of the temporal evolution of $f_{\\mathrm{dust}}$ and $R_{\\mathrm{dust}}$ based on these values together with previously estimated stellar ages, and find that our results are consistent with a $f_{\\mathrm{dust}}$\\,$\\propto$\\,$t^{-\\alpha}$, with $\\alpha$\\,$\\sim$\\,0.8--2.0, as well as tentatively with the $R_{\\mathrm{dust}}$\\,$\\propto$\\,$t^{1/3}$ relation predicted in debris disks collisional evolution models by \\citet{Kenyon2004a}. \\end{itemize}"
    },
    "1005/1005.3165_arXiv.txt": {
        "abstract": "{} {Gigahertz peaked spectrum (GPS) radio sources are compact extragalactic radio sources, assumed to be young radio-loud active galactic nuclei and ideal objects for studying the early evolution of extragalactic radio sources. The Very Long Baseline Interferometry (VLBI) observation with high resolution is vital for studying the structure of these compact sources. } {We defined a sub-sample of twelve GPS sources which have not been observed with the VLBI before, from the Parkes half-Jansky sample, and carried out VLBI observations at 1.6 GHz and 5 GHz with the European VLBI Network (EVN) in 2006 and 2008, respectively, to classify the source structure and to find compact symmetric objects (CSOs). Additionally, we carried out the 4.85 GHz flux density observations for these sources with the Urumqi 25-m telescope between the years 2007 and 2009 to study whether there is any variability in the total flux density of the GPS sources.} {The results of the 5 GHz VLBI observations and total flux densities of these sources are presented in this paper. From the VLBI morphologies, the spectral indices of components and the total flux variability of the twelve targets, we firmly classify three sources J0210+0419, J1135$-$0021, and J2058+0540 as CSOs, and classify J1057+0012, J1203+0414, and J1600$-$0037 as core-jet sources. The others J0323+0534, J0433$-$0229, J0913+1454, J1109+1043, and J1352+0232 are labelled CSO candidates, and J1352+1107 is a complex feature. Apart from core-jet sources, the total flux densities of the CSOs and candidates are quite stable at 5 GHz both during a long-term of $\\sim$20 years relative to the PKS90 data and in a period between 2007 and 2009. The total flux densities are resolved-out by more than 20\\% in the 5 GHz VLBI images for 6 sources, probably because of diffuse emission. In addition, we estimated the jet viewing angles ($\\Theta$) for the confirmed CSOs by using the double-lobe flux ratio of the sources, the result being indicative of relatively large $\\Theta$ for the CSOs.} {} ",
        "introduction": "Gigahertz peaked spectrum (GPS) radio sources are powerful ($P_{\\rm1.4~GHz}\\geq\\rm10^{25}\\,W\\,Hz^{-1}$) and compact extragalactic radio sources characterized by a convex radio spectrum peaking at between 0.3 and a few GHz (see O'Dea 1998 for a review). Sources that peak at $>$ 5 GHz are called high frequency peakers (HFP, Dallacasa et al. 2000), which are understood to be even more compact and younger. The spectrum turnover of the GPS sources is caused by synchrotron self-absorption (de Vries et al. 2009), in general, and it is also possible that free-free absorption plays a role in producing the spectral turnover in some cases (e.g., Luo et al. 2007; Orienti \\& Dallacasa 2008).   Extended emission in GPS source is uncommon especially for GPS galaxies, most of which are truly compact ($< 1\\,kpc\\,h^{-1}$) (Stanghellini et al. 2005). The compactness of GPS sources is most likely due to their youth ($< 10^4$ years) according to a spectral ageing analysis (Murgia et al. 2003). The youth scenario is supported by the hotspot kinematics in a couple of GPS sources that are also identified as compact symmetric objects (CSOs) (Giroletti \\& Polatidis 2009). A unification model assumes that GPSs/CSOs will evolve into compact steep spectrum (CSS) sources or medium-sized symmetric objects (MSOs, $1-20\\,kpc\\,h^{-1}$), which in turn, evolve into large radio galaxies ($> 20\\,kpc\\,h^{-1}$), i.e., FRI/II radio sources (Fanti et al. 1995, 2009; Snellen et al. 2000).  Furthermore, the GPS sources are associated with either galaxies or quasars. The GPS quasars usually exhibit some form of core-jet structure, while the GPS galaxies often show CSO-like or compact doubles (CD)(e.g., Orienti et al. 2006, Liu et al. 2007), suggesting that the GPS galaxies have relatively large jet-viewing angles. The GPS galaxies have very low polarization (less than 0.5\\% at 5~GHz, Xiang et al. 2006), which is probably because of the large Faraday depth in the line of sight, implying that their host AGN are largely edge-on to us.  Complete samples of GPS sources are statistically vital for studying the early evolution of radio-loud AGN and the triggering of their nuclear activity, such as the 1 Jy complete GPS sample in the northern sky (Stanghellini et al. 1998), the Parkes half-Jansky (PHJ) sample in the southern/equatorial sky (Snellen et al. 2002), and the CORALZ sample at low redshift (Snellen et al. 2004). The Very Long Baseline Interferometry (VLBI) technique is a key tool for studying the morphologies of GPS sources from the GPS samples. We defined a subsample of twelve GPS sources with declination $>-5^{\\circ}$ that had not been observed with the VLBI before, from the PHJ sample, and carried out the VLBI observations at 1.6 GHz and 5 GHz of them in 2006 and 2008, respectively. It is found that most of the sources exhibit compact double structure by analysing the 1.6 GHz European VLBI Network (EVN) observation (Liu et al. 2007, hereafter paper I). In the follow-up 5 GHz EVN observation, we demonstrate that we can resolve their VLBI structure more sharply and obtain the spectral indices of components between 1.6 GHz and 5 GHz, and finally identify CSOs from the candidates.  Quite a few GPS sources having been found to have flux variability (Torniainen et al. 2007). Since 2007, we have performed flux monitoring at 4.85 GHz with the Urumqi 25-m radio telescope for GPS sources from the master list for which declination $>-30^{\\circ}$ (Labiano et al 2007).  Throughout the paper, the cosmological parameters $H_{0}=70.1$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.721$ and $\\Omega_{m}=0.279$ (Komatsu et al. 2009) are adopted, and the spectral index $\\alpha$ is defined by $S\\propto\\nu^{-\\alpha}$. ",
        "conclusions": "We have presented the results of 5 GHz EVN observations of 12 GPS sources from the sample of Snellen (2002) for the first time. The source structure and spectral indices between 1.6 GHz and 5 GHz are obtained. We have also carried out total flux density measurements with the Urumqi 25-m telescope at 4.85 GHz in 2007--2009, and compared the mean flux density with the PKS90 data.  \\begin{enumerate}  \\item From the source morphologies, the component spectral indices and the total flux variability, we can firmly classify three sources J0210+0419, J1135$-$0021, and J2058+0540 as CSOs, and J1057+0012, J1203+0414, and J1600$-$0037 as core-jet sources. The others J0323+0534, J0433$-$0229, J0913+1454, J1109+1043, and J1352+0232 remain as CSO candidates, and J1352+1107 is a complex VLBI feature whose nature need to clarify in the future.  \\item The total flux densities of the GPS sources at 5 GHz were quite stable in 2007--2009 and over the long term compared with the PKS90 data, except for core-jet sources.  \\item We estimated the jet viewing angles $\\Theta$ for the CSOs by using the double-lobe flux ratio. The three CSOs were found to have systematically larger $\\Theta$.  \\item We plan to observe the CSOs in the future with the VLBI at 5 GHz to measure the expansion speeds of lobes and estimate the ages of these young radio sources.  \\end{enumerate}"
    },
    "1005/1005.3492_arXiv.txt": {
        "abstract": "Recent studies have shown that the primordial non-Gaussianity affects clustering of dark matter halos through a scale-dependent bias and various constraints on the non-Gaussianity through this scale-dependent bias have been placed.  Here we introduce the cross-correlation between the CMB lensing potential and the galaxy angular distribution to effectively extract information about the bias from the galaxy distribution.  Then, we estimate the error of non-linear parameter, $f_{\\rm NL}$, for the on-going CMB experiments and galaxy surveys, such as Planck and Hyper Suprime-Cam (HSC). We found that for the constraint on $f_{\\rm NL}$ with Planck and HSC, the wide field galaxy survey is preferable to the deep one, and the expected error on $f_{\\rm NL}$ can be as small as: $\\Delta f_{\\rm NL} \\sim 20$ for $b_0 = 2$ and $\\Delta f_{\\rm NL} \\sim 10$ for $b_0 = 4$, where $b_0$ is the linear bias parameter.  It is also found that future wide field galaxy survey could achieve $\\Delta f_{\\rm NL} \\sim 5$ with CMB prior from Planck if one could observe highly biased objects at higher redshift ($z\\sim 2$). ",
        "introduction": "\\label{sec:Intro} The cosmic microwave background (CMB) temperature anisotropy is a quite useful probe for cosmology. The contribution to the anisotropy is dominated by fluctuations at the last scattering surface. The CMB photon, however, encounters the large-scale structure along the line of sight and some additional effects are imprinted on the temperature and polarization as secondary anisotropies. The deflection of the CMB photon due to gravitational potential produced by the large-scale structure is one of them. The effect of the gravitational lensing on the CMB photon through the large-scale structure is known as the CMB lensing \\cite{Lewis-06}. An on-going CMB observation by Planck \\cite{Planck} or various ground-based experiments are expected to detect this signal, while the effect of the CMB lensing are imprinted on small scales which the Wilkinson Microwave Anisotropy Probe (WMAP) satellite could not resolve. The lensing effect reflects the late time evolution of the universe at relatively low-redshifts. Therefore, the lensing information plays an important role to determine the cosmological parameters, such as the neutrino mass, the cosmological constant, the equation of state parameter of dark energy and so on.  The large-scale structures are formed at relatively late time and they become the source of the gravitational potential. They are correlated with the CMB temperature anisotropy through the Integrated Sachs-Wolfe (ISW) effect, which generates the secondary anisotropies due to the time variation of the potential \\cite{Sachs-67}. The cross-correlation has an advantage for observations of ISW effect whose signal is weak. Cross-correlations with complementary probes are expected to provide additional information on top of their respective auto-correlations. Similarly, we expect that lensing potential should correlate with the large-scale structure and the information from their cross-correlation may precisely determine the cosmological parameters. \\\\  Recently, the deviations from Gaussian initial conditions (primordial non-Gaussianity) are  intensively focused on and discussed. Inspection of them offers an important window into the very early Universe because non-standard models of inflation allow for a large non-Gaussianity while standard single-field slow-roll models predict the small deviations from Gaussianity. The most popular method to detect the primordial non-Gaussianity is to measure higher-order correlations of CMB anisotropies and distributions of galaxies, for example, the bispectrum or the three-point correlation function of CMB \\cite{Komatsu-01, Bartolo-04} or the large-scale structure bispectrum \\cite{Scoccimarro-00, Verde-00, Scoccimarro-04, Sefusatti-07}.  Some studies have shown that the primordial non-Gaussianity affects clustering of dark matter halos through a scale-dependent bias, both by analytic calculations and by N-body simulations \\cite{Afshordi-08, Carbone-08, Dalal-08, Pillepich-10, Slosar-08}.  By considering the scale-dependent bias, one can constrain on the non-Gaussianity by the power spectrum. Various constraints on the non-Gaussianity through the scale-dependent bias have already been placed \\cite{Slosar-08, Oguri-09, Carbone-10, Cunha-10, DeBernardis-10, Xia-10}. One may expect, however, that a constraint on non-linear parameter $f_{\\rm NL}$, which describes the primordial non-Gaussianity, is degenerated with other cosmological parameters and has a large error.  Because of such degeneracy, especially with the linear bias $b_0$, it is important to combine unbiased observations which are sensitive to the matter power spectrum in the near universe, like CMB lensing, shear and so on.  Here we use the cross-correlation between the CMB lensing potential and the large-scale galaxy distribution and estimate expected errors of non-linear parameter $f_{\\rm NL}$ for future CMB experiments and galaxy surveys.  We expect that this cross-correlation may play a role to break degeneracy and give us more stringent constraints.  This paper is organized as follows. We review the scale-dependent bias due to the non-Gaussianity in section \\ref{sec:SBias} and the theory behind the cross-correlation between CMB lensing potential and galaxy distribution in section \\ref{sec:APS}. In section \\ref{sec:Modeling} we describe the survey model for the galaxy distribution in the Hyper Suprime-Cam (HSC) survey, which is a fully funded imaging survey at Subaru telescope. In section \\ref{sec:Fisher} and section \\ref{sec:Forecasts} we explain the method of our analysis. Finally, in section \\ref{sec:Result} and \\ref{sec:Summary} we discuss the results and summarize our conclusions. Throughout this paper we assume a spatially flat universe for simplicity. ",
        "conclusions": "\\label{sec:Summary} In this paper we have estimated the constraints on the cosmological parameters newly taking into account the cross-correlation between CMB lensing and galaxy angular distributions. In particular, we have focused on the constraint on the primordial non-Gaussianity through the scale-dependent bias $b (z, k)$ and estimated how much the cross-correlation between CMB lensing and galaxy would improve the constraint by the Fisher matrix analysis. In order to make the most general Fisher matrix analysis with CMB and galaxy survey experiments, we have taken into account the all the auto- and cross-correlations available which is expressed by the 8$\\times$8 covariance matrix as Eq. (\\ref{eq:Cov}). We have paid particular attention to the CMB and galaxy survey just coming up now, namely Planck satellite and HSC, and also to the future experiments like CMBPol, LSST and so on. Our estimations are mainly based on the Planck and HSC surveys, however we also show some cases for comparison when CMBPol or ambitious survey conditions are assumed.   As for the constraints on the conventional cosmological parameters, the improvement can not be expected very much from the simple estimate in which the Fisher matrices for the CMB and galaxy surveys are combined properly even if the cross-correlations are properly taken into account.  However, focusing on the determination of $f_{\\rm NL}$, we found that the role of the cross-correlation between CMB and galaxy is important, especially the one between CMB lensing potential and galaxy $\\cpg$ contributes the determination of $f_{\\rm NL}$.   We have estimated the constraints on the $f_{\\rm NL}$ for the coming experiments.  First, we gave the rough estimation in cases where the survey area $f_{\\rm sky}$ and galaxy samples $N_{\\rm g}$ are fixed and only peak redshift $z_0$ is varied. It was found that the keys for more strictly constraining the primordial non-Gaussianity are observing higher redshift and larger biased objects. Second, considering the realistic observations with fixed observation time, we have estimated the constraint on $f_{\\rm NL}$, adapting the galaxy survey model in Ref.~\\cite{Yamamoto-07}, which is scaled to the HSC survey. As a result we found on optimized target redshift to be $z_0 \\simeq 1$, which brings us to a conclusion that we should make a wide survey rather than a deep survey for constraining the primordial non-Gaussianity with HSC-like observation. This is because the effect of the primordial non-Gaussianity through the scale-dependent bias is significant on large scales rather than small scales. The large-scale region is dominated by the cosmic variance related to the survey area while the small-scale region is dominated by the shot noise term related to the number of the galaxy samples. Therefore, we should explore the wide survey area rather than observing a lot of galaxies. However, we should keep it in mind that high-redshift and highly-biased objects are much affected by primordial non-Gaussianity, so that the deep survey will be also essential, in the future.   The constraints on the primordial non-Gaussianity expected from HSC survey with Planck are : $\\Delta f_{\\rm NL} \\sim ~20$ for $b_0 = 2$ and $\\Delta f_{\\rm NL} \\sim ~10$ for $b_0 = 4$.  Slosar et al.~\\cite{Slosar-08} obtained constraints on $f_{\\rm NL}$ for highly biased tracers using available luminous red galaxy (LRG) or quasar (QSO) data from Sloan Digital Sky Survey (SDSS) \\cite{Ho-08} and CMB data from WMAP 5, with the error of $\\Delta f_{\\rm NL} \\sim 30$.  The reason why the combination of HSC and Planck observations does not make significant improvement over the current constraints is explained as below.  We have considered only 800 hours for HSC and normal galaxies. The SDSS data is obtained over longer period of time, and QSOs seem to be observed more than normal galaxies at high-redshift. These constraints are weaker than those expected with the CMB bispectrum constraints achievable with an ideal CMB experiment, $\\Delta f_{\\rm NL} \\sim ~{\\rm few}$ (Ref.~\\cite{Yadav-07, Liguori-08}).  However, the constraint on $f_{\\rm NL}$ presented in this paper depends highly on the galaxy survey condition, $i.e.,$ survey area $f_{\\rm sky}$, number density of sample galaxies $n_{\\rm L}$ and observing peak redshift $z_0$.  It is found that with some galaxy survey with Planck $f_{\\rm NL}$ could achieve $\\Delta f_{\\rm NL} \\sim ~ 5$ at ${z_{\\rm m}} = 2.0$ (Fig.~\\ref{fig:dfNL-FS}). Recently it is reported that an ambitious future galaxy survey (like the LSST survey), which provides large survey area of 30,000 ${\\rm deg^2}$ and highly biased galaxy samples, can measure the primordial non-Gaussianity with the order $\\Delta f_{\\rm NL} \\sim ~2-5$ \\cite{Carbone-10}.  The other method using the full covariance of cluster counts for Dark Energy Survey (DES) can yield $\\Delta f_{\\rm NL} \\sim 1-5$ \\cite{Cunha-10}.  In Ref.~\\cite{Carbone-10, Cunha-10}, they add the Planck and galaxy survey Fisher matrices, $i.e.,$ $F_{ij} = F_{ij}^{\\rm galaxy} + F_{ij}^{\\rm CMB} $, and do not include the cross-correlation between CMB and galaxy survey, $i.e.,$ $\\ctg$ an $\\cpg$, so that more tight constraint on $f_{\\rm NL}$ may be expected with these cross-correlations.  In any case, it is worth pursuing how well we can put a constraint on non-Gaussianity of the local-type from the large-scale structure because it contains information on non-Gaussianity at different epoch from CMB and thus the constraint through the scale-dependent bias will be an important cross check against the CMB bispectrum.   \\begin{figure}[t] \\centering \\includegraphics[clip,keepaspectratio=true,width=0.45\\textwidth]{fig/dfNL-0530-40.eps} \\hspace{5mm} \\includegraphics[clip,keepaspectratio=true,width=0.45\\textwidth]{fig/fnl-HSC_0530.eps} \\caption{ The 1$\\sigma$ error of $f_{\\rm NL}$ in the case with fixed sky coverage $f_{\\rm sky} = 0.1$ and total number of galaxy $N_{\\rm g} = 10^6$ (Left), and in the case with modeling for HSC survey (Right). For the CMB experiment we consider two cases. The thick lines show the errors against peak redshift $z_0$ and the thin lines are against mean redshift $z_{\\rm m}$. \\label{fig:ParErrorM}} \\end{figure} \\begin{table}[th] \\begin{ruledtabular} \\begin{tabular}{ccccc} $z_0$ & $z_{\\rm m}$ & $\\Delta f_{\\rm NL}$ & $n_{\\rm L} {\\rm [arcmin^{-2}]}$ & $f_{\\rm sky}$\\\\ \\hline 0.8 & 0.51 & 30.1 &  0.9 & 0.42 \\\\ 0.9 & 0.58 & 24.5 &  1.9 & 0.41 \\\\ 1.0 & 0.64 & 20.7 &  3.7 & 0.40 \\\\ 1.1 & 0.70 & 18.5 &  7.0 & 0.35 \\\\ 1.2 & 0.77 & 18.6 & 12.4 & 0.25  \\\\ 1.3 & 0.83 & 21.9 & 20.9 & 0.14 \\\\ 1.4 & 0.90 & 29.1 & 34.0 & 0.06 \\\\ 1.5 & 0.96 & 40.5 & 53.5 & 0.02 \\end{tabular} \\caption{ The relation of parameters with fixed observation time, between the peak redshift $z_0$, the mean redshift $z_{\\rm m}$, the $1\\sigma$ error of $f_{\\rm NL}$, the number density of sampling galaxy $n_{\\rm L}$ and the survey area $f_{\\rm sky}$, for the HSC-like survey ($b_0=2$). \\label{tb:ParRelation}} \\end{ruledtabular} \\end{table}"
    },
    "1005/1005.1204_arXiv.txt": {
        "abstract": "{Seven years of successful observations of the sky have been completed within the INTEGRAL mission, in the transition regime between X-rays and $\\gamma$-rays from $\\sim$10-8000~keV. Initially-agreed mission goals have been pursued, and both high-resolution images of point sources and high-resolution spectra of nuclear lines have been obtained. New discoveries have been made, such as X-ray emission from embedded binaries, hard emission tails from AXPs and $^{60}$Fe radioactivity lines; these stimulated both theoretical and observational studies, and now make INTEGRAL a valuable asset for the astronomical survey of high-energy sources across the sky. This contribution summarizes the situation after seven years of the mission, and concludes the 7-year anniversary workshop {\\it The extreme sky} held in Otranto, Italy, in Oct 2009.}  \\FullConference{INTEGRAL Workshop\\\\ {\\it The Extreme Universe} \\\\ October 13-17, 2009\\\\ Otranto, Italy}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.2031_arXiv.txt": {
        "abstract": "{} {`Water in Star-forming regions with \\textit{Herschel}' (WISH) is a \\textit{Herschel} Key Programme aimed at understanding the physical and chemical structure of young stellar objects (YSOs) with a focus on water and related species.} {The low-mass protostar HH~46 was observed with the Photodetector Array Camera and Spectrometer (PACS) on  the \\textit{Herschel} Space Observatory to measure emission in H$_2$O, CO, OH, [\\ion{O}{i}], and [\\ion{C}{ii}] lines located between 63 and 186 $\\mu$m. The excitation and spatial distribution of emission can disentangle the different heating mechanisms of YSOs, with better spatial resolution and sensitivity than previously possible.} {Far-IR line emission is detected at the position of the protostar and along the outflow axis. The OH emission is concentrated at the central position, CO emission is bright at the central position and along the outflow, and H$_2$O emission is concentrated in the outflow. In addition, [\\ion{O}{i}] emission is seen in low-velocity gas, assumed to be related to the envelope, and is also seen shifted up to 170 km~s$^{-1}$ in both the red- and blue-shifted jets. Envelope models are constructed based on previous observational constraints. They indicate that passive heating of a spherical envelope by the protostellar luminosity cannot explain the high-excitation molecular gas detected with PACS, including CO lines with upper levels at $>$2500 K above the ground state. Instead, warm CO and H$_2$O emission is probably produced in the walls of an outflow-carved cavity in the envelope, which are heated by UV photons and non-dissociative $C$-type shocks. The bright OH and [\\ion{O}{i}] emission is attributed to $J$-type shocks in dense gas close to the protostar. In the scenario described here, the combined cooling by far-IR lines within the central spatial pixel is estimated to be 2$\\times$10$^{-2}$ $L_\\odot$, with 60--80\\% attributed to $J$- and $C$-type shocks produced by interactions between the jet and the envelope.} {} ",
        "introduction": "The embedded phase of star formation is a critical period in the evolution of a young star, because it is the stage where the final mass of the star, the size and mass of the protoplanetary disk, and the initial chemical composition of the disk are determined \\citep{Andre00,Visser09b,Jorgensen09}. Many physical processes occur simultaneously in the immediate surroundings of the protostar: infall in the collapsing envelope, outflows sweeping up and shocking the material, and energetic (UV and X-ray) photons heating and dissociating the gas \\citep{Spaans95,Bachiller99,Arce07}. Because of high extinction, these processes can only be probed at far-infrared and millimetre wavelengths, but lack of observational facilities has hampered their quantification.  The goal of the `Water In Star-forming regions with {\\it Herschel}' (WISH) key programme is to use H$_2$O, CO and related species to determine the physical and chemical characteristics of young stellar objects (YSOs) as functions of evolutionary stage and across a wide range of luminosities and masses \\citep[see also][]{Nisini10, Fich10}.  One of the many surprises of the Infrared Space Observatory (ISO) was the detection of highly excited CO and H$_2$O lines towards low-mass YSOs with the Long Wavelength Spectrometer \\citep[LWS; e.g.,][]{Giannini99,Ceccarelli99, Nisini02}. The origin of this hot gas ($\\sim$500--2000 K) has been heavily debated, with two main explanations put forward: (i) shocks extending over a large area (arcmin scale), and (ii) the envelope heated by the protostellar luminosity on scales of $2-20$\\arcsec. The ISO-LWS data and subsequent modelling could not distinguish between these scenarios given the large beam of $\\sim$80$''$. The smaller aperture of {\\it Herschel}-PACS allows imaging of these lines at 9\\farcs4, sufficient to separate the on- and off-source emission. Moreover, its higher sensitivity and higher spectral resolution allow detection of weaker lines.  This letter presents Science Demonstration Phase observations of HH~46, an isolated low-mass protostar ($L_{\\rm bol}\\approx16\\ L_\\odot$, $D\\approx450$ pc) located in a dense core with a prominent outflow extending out to $\\sim$2$'$. It has been imaged with the {\\it Spitzer} Space Telescope \\citep{Velusamy07} and ground-based sub-millimetre telescopes \\citep[and references therein]{vanKempen09a}. Because of its well-defined geometry, HH~46 provides an ideal testbed for separating the different physical components contributing to the observed emission and for benchmarking new models of YSOs that can be used for other sources. ",
        "conclusions": ""
    },
    "1005/1005.5611_arXiv.txt": {
        "abstract": "In this work we propose a new numerical approach to distinguish between regular and chaotic orbits in Hamiltonian systems, based on the simultaneous integration of both the orbit and the deviation vectors using a symplectic scheme, hereby called {\\em global symplectic integrator}. In particular, the proposed method allows us to recover the correct orbits character with very large integration time steps, small energy losses and short CPU times. To illustrate the numerical performances of the global symplectic integrator we will apply it to two well-known and widely studied problems: the H\\'enon-Heiles model and the restricted three-body problem. ",
        "introduction": "\\label{sec:intro}   Hamiltonian systems exhibit phase spaces where regular and chaotic orbits do coexist. This fact makes the problem of the numerical characterization of regular and chaotic motion an hard task, notably in systems with many degrees of freedom. For this reason scientists developed fast and accurate tools to obtain information about the chaotic versus regular nature of the orbits of such systems, and efficiently characterize large domains in their phase space as ordered or (weakly) chaotic.  These methods can be roughly divided into two major groups: {\\em Lyapunov--like} methods, i.e. methods based on the study of the evolution of deviation vectors for a given orbit; the methods FLI~\\cite{Froe1997}, \\texttt{MEGNO}~\\cite{cincotta2000}, \\texttt{SALI}~\\cite{S01} and \\texttt{GALI}~\\cite{SBA07} belong to this first class. And {\\em Fourier--like} methods based on the determination of the frequencies of the spectrum of some observable related to a given orbit, for instance \\texttt{FMA}~\\cite{Laskar1992,Laskar1993} or the Spectral method~\\cite{GB2001}. In the present work we will be interested in the former class, in particular to the \\texttt{SALI} chaos detection technique.  To numerically compute such an indicator, one needs to have a good orbit determination, but also to integrate the evolution of two deviation vectors. It is well-known that symplectic integration schemes outperform the non--symplectic ones, when compared using the same order of accuracy and the same integration step size. This is mainly due to the very good energy conservation properties, but also to other first integrals. As a consequence, larger step sizes are allowed while still keeping a reasonably small energy loss. Thus they allow us to enlarge considerably the time span of the numerical simulation, without degradating the goodness of the numerical results, that is why symplectic integrators turn out to be essential for the long-term evolution of a dynamical system.  In the same way, chaos detection techniques could also benefit from the use of symplectic integrators. The aim of this paper is to show that the symplectic integration of the deviation vectors can improve the ability of the previous chaos indicators in the characterization of regular and chaotic orbits by correctly identifying the orbit behavior, using a larger integration step size. Because our method propose to integrate simultaneously both the orbit and the deviation vectors, using a symplectic scheme, we hereby name this method {\\em global symplectic integrator}.  More precisely, we will accurately determine, using a symplectic scheme and a small enough integration time step $\\tau$, the character, i.e. the regular or chaotic behavior, of a large number, $N_{tot}$, of orbits using one of the aforementioned chaos detectors. Then, all the previous orbits will be reanalyzed using both symplectic and non-symplectic integrators schemes, but with larger and larger time steps. We will show that the use of the global symplectic integrator will allow us to recover a large percentage of orbits' characters with very large time steps compared to the non--symplectic one and using small amount of CPU time.  Our symplectic method will be applied to the classical H\\'enon--Heiles model, and also to a well-known problem of celestial mechanics: the restricted three-body problem, where the possibility of using large integration step sizes is essential for the study of secular resonances.  The paper is organized as follows. In Section~\\ref{sec:themethod}, we describe our method for integrating symplectically the deviation vectors used in the chaos detection techniques. For the sake of completeness, we present, in Section~\\ref{sec:sali}, a brief introduction to the \\texttt{SALI} chaos indicator. The results of our method applied to the H\\'enon--Heiles system are presented in Section~\\ref{ssec:hh}, while Section~\\ref{ssec:3b} will be devoted to present an application to the restricted three-body problem. Finally in Section~\\ref{sec:ccl} we will summarize our findings and we will conclude. ",
        "conclusions": "\\label{sec:ccl}  In this work, we proposed a new method for the detection of regular and chaotic orbits in Hamiltonian systems, based on the integration of the deviation vectors used in chaos detection techniques, using symplectic algorithms. Our method has been tested on two well-known models, and the results clearly demonstrate that it outperforms non-symplectic ones.  Concerning the H\\'enon-Heiles system, it appears that, for large time steps, non-symplectic integrators tend to detect an excessive number of chaotic orbits, while the global symplectic integrator is able to identify correctly the character of nearly all orbits for larger time steps, up to four times larger than non--symplectic ones. Moreover, due to his symplectic properties, we showed that our method ensures a very small energy loss even on very long time spans. Let us emphasize that the possibility to use larger time steps saves a considerable amount of computation time.  This possibility turns out to be essential for the study of the Kozai resonance in the restricted three-body problem, where the secular orbital changes operate on extremely long time scales. Once again, the influence of the separatrix of this problem is well identified by our global symplectic integrator. On the contrary, the Bulirsch-Stoer non-symplectic method, even with smaller time steps, seems unable to distinguish between regular and chaotic motion, at least on the same integration time span. A possible reason for this behavior could be the accumulation of numerical errors introduced by the integrator and a significant energy loss, disadvantages which are avoided using our symplectic scheme.  We are confident that our findings should be generic for a large class of Hamiltonian systems. Thus we encourage scientists working on chaos indicators to perform symplectic integrations of both the orbit and the deviation vectors using the global symplectic integrator, as proposed in the present paper, whenever the Hamiltonian is of the form $H(\\vec{x})=A(\\vec{p})+B(\\vec{q})$, or generically it can be divided into two parts, each one separately integrable. Computation time and reliability of the results could thus benefit a lot from this procedure, as we clearly demonstrated above."
    },
    "1005/1005.0491_arXiv.txt": {
        "abstract": "{ \\textsc{Context}. One of the major challenges in modern astrophysics is to understand the origin and the evolution of galaxies, the bright, massive early type galaxies (ETGs) in particular. There is strong observational evidence that massive ETGs are already in place at redshift $z \\sim 2-3$ and that they formed most of their stars well before $z=1$. Therefore, these galaxies are likely to be good probes of galaxy evolution, star formation and, metal enrichment in the early Universe.  \\noindent \\textsc{Aims}. In this context it is very important to set up a diagnostic tool able to combine results from chemo-dynamical N-Body-TSPH (NB-TSPH) simulations of ETGs with those of spectro-photometric population synthesis and evolution so that all key properties of galaxies can be investigated. These go from the integrated spectrum and magnitudes in any photometry, both in the rest-frame and as a function of the redshift, to present-day structural properties. The main goal of this paper is to provide a preliminary validation of the software package before applying it to the analysis of observational data.  \\noindent \\textsc{Methods}. The galaxy models in use where calculated by the Padova group in two different cosmological scenarios: the standard cold dark matter cosmology ($SCDM$), and the so-called Concordance cosmology ($\\Lambda CDM$, with $\\Omega_\\Lambda=0.762$). For these template galaxies, we recover their spectro-photometric evolution through the entire history of the Universe. This is done in particular for two important photometric systems, the Bessell-Brett and the Sloan Digital Sky Survey (SDSS) passbands.  \\noindent \\textsc{Results}. We computed magnitudes and colors and their evolution with the redshift along with the evolutionary and cosmological corrections for the model galaxies at our disposal, and compared them with data for ETGs taken from the COSMOS and the GOODS databases. Finally, starting from the dynamical simulations and photometric models at our disposal, we created synthetic images in a given photometric system, from which we derived the structural and morphological parameters. In addition to this, we address the question of the scaling relations, and in particular we examine the one by Kormendy. The theoretical results are compared with observational data of ETGs selected form the SDSS database.  \\noindent \\textsc{Conclusions}. The simulated colors for the different cosmological scenarios follow the general trend shown by galaxies of the COSMOS and GOODS surveys at lower redshifts and are in good agreement with the data up to $z \\sim 1$, where the number of early-type galaxies observed falls abruptly. In conclusion, within the redshift range considered, all the simulated colors reproduce the observational data quite well. Looking at the structural parameters derived from the surface imaging, the luminosities and effective radii (Kormendy relation) measured for our model galaxies are consistent with the archival data from the SDSS.} ",
        "introduction": "The origin and evolution of early-type galaxies (ETGs), the bright massive ETGs in particular, are two of the major challenges in modern astrophysics, and it is still a very controversial subject \\citep{Chiosi00}. Spheroidal systems are of interest in their own right because they contain more than half of the total stellar mass in the local Universe \\citep{Fukugita98}. Giant ETGs appear to define a homogeneous class of objects that predominantly consists of uniformly old and red populations, which implies that they must have formed at high redshift and that they have negligible amounts of gas and very little star formation \\citep{Bressan94}.  There is strong observational evidence that old, massive, red, and metal-rich proto-ETGs are already in place at $z \\sim$2-3 and that the present-day early-type galaxies formed most of their stars well before redshift $z=1$ \\citep{Searle73,Brinchmann00,Treu05,vanderWel05}. Moreover, the current rates of star formation in these systems are quite low, whereas the rates increase sharply into the past \\citep{Butcher78,Dressler80}. Therefore, these ETGs are likely good probes of galaxy assembly, star formation, and metal enrichment in the early Universe.  \\noindent \\textsf{The cosmological background.} In a Universe dominated by cold dark matter (CDM), some kind of dark energy in form of the cosmological constant $\\Lambda$, and containing a suitable mix of baryons and photons, the cosmic structures are formed by the gravitational collapse of dark matter and are organized in a hierarchy of halos inside which baryons dissipate their energy and collapse to form luminous systems. The formation of ETGs can be reduced to the following schemes \\citep{Peebles02,Schade99}:  \\noindent i) {\\it Early, monolithic-like aggregation}. This scenario of galaxy formation predicts that all ETGs form at high redshift ($z \\gg 1$) as a result of rapid and dissipation-less collapse of a large mass of gas soon transformed into stars. In the model, first proposed by \\citet{Eggen62} and then refined and improved by \\citet{Larson75},\\citet{Arimoto87},\\citet{Bressan94},and \\citet{Chiocar02}, ETGs undergo a single and short, but intense, burst of star formation, followed ever since by the passive evolution of their stellar populations to the present day. This simple model naturally accounts for the old ages ($\\sim 12$\\,Gyr) of spheroidal galaxies, their high densities, and the weak temporal evolution of their stellar content.  In favor of this scheme are the observational data that convincingly hint at old and homogeneous stellar populations \\citep[see][ for a review of the subject]{Chiosi00}. It is worth mentioning, however, that \\citet{Kauffmann93} and \\citet{Barger99} argue for some recent evolution in the stellar populations of elliptical galaxies. This scenario reproduces the optical properties of ETGs remarkably well, and successfully explains the tightness of the fundamental scaling relations that ETGs obey, like the color-magnitude relation and the fundamental plane, as well as the evolution of these relations as a function of redshift.  The monolithic formation mechanism fails to explain some recent observational evidence that has become available with the advance of more detailed data from present-day surveys. These indicate that the star formation histories of at least some ETGs, and perhaps the early-type population as a whole, deviate strongly from the expectations of the monolithic collapse paradigm, both in terms of their structural evolution and star formation experienced by them over the whole Hubble time. It is less successful at explaining the detailed luminosity dependence of their dynamical properties, the apparent scarcity of very large star-bursts in the high-redshift universe, and the origin of dynamical peculiarities indicating some recent accretion events. This scenario does not fit the currently accepted $\\Lambda CDM$ picture of galaxy formation whose bottom line is that massive dark matter halos are assembled by mergers of low-mass halos and therefore the mass of a galaxy is thought to accumulate over the lifetime of the Universe.  To cope with some of the above difficulties, a hybrid scenario named {\\it revised monolithic} has been proposed by \\citet{Schade99} and confirmed by NB-TSPH simulations by \\citet{Merlin06,Merlin07}, who suggest that a large number of the stars in massive galaxies are formed very early-on at high redshift ($z \\sim$1-2) and the remaining few at lower $z$. The revised monolithic ought to be preferred to the classical monolithic, as some evidence of star formation at 0.2$\\leq z\\leq$2 can be inferred from the emission line of [OII], and also the number frequency of ETGs up to $z\\simeq$1 seems to be nearly constant. Recently, \\citet{PerezGonzalez08}, whom analyzing a huge sample of galaxies, confirmed a scenario where most massive objects assemble their mass very early, whereas the smallest galaxies evolve more slowly building up their mass at lower redshift.  \\noindent ii) {\\it The hierarchical aggregation.} This scenario instead suggests that massive ETGs are the end product of subsequent violent mergers of preexisting smaller subunits, on time scales almost equal to the Hubble.  In this scenario, the epoch of assembly of ETGs differs markedly from the epoch of formation of their constituent stars, and the high density of elliptical galaxies is ascribed to the effects of dissipation during the formation of the progenitor disks. As the look-back time increases, the density in comoving space of bright (massive) ETGs should decrease by a factor 2 to 3 \\citep[see e.g.][]{White78,Kauffmann93}.  This model accounts naturally for the scarcity of very bright elliptical progenitors at high redshift, for the rapid evolution of the galaxy population with look-back time, and for dynamical peculiarities. In favor of this view is some observational evidence that the merger rate likely increases with $\\sim(1+z)^{3}$ \\citep{Patton97}, together with some hint for a color-structure relationship for E \\& S0 galaxies: the color becomes bluer at increasing complexity of a galaxy structure. This could indicate some star formation associated to the merger event. Finally, there are the many successful numerical simulations of galaxy encounters, mergers, and interactions \\citep[e.g.][]{Barneshern96}.  It is, on the other hand, less successful in explaining the apparent old ages of stars in elliptical galaxies and their uniformity in dynamical properties. Nevertheless, contrary to the expectation from this model, the number density of ellipticals do not seem to decrease with the redshift, at least up to $z\\simeq 1$ \\citep{Im96}. A significant population of massive and passive ETGs up to $z \\approx 2.5$ and some hints about massive ETGs at redshift $z>3$ \\citep{Cimatti09}, clearly do not agree with the classical hierarchical scenario, because we need to have big objects already in place at higher and higher $z$.  There is a companion scheme named \\textit{dry merger}, in which bright ETGs form by encounters of quiescent, no star-forming galaxies. This view is advocated by \\citet{Bell04}, who find that the B-band luminosity density of the red peak in the color distribution of galaxies shows mild evolution starting from $z\\simeq$1. As old stellar populations would fade by a factor 2 or 3 in this time interval, and the red color of the peak tells us that new stars are not being formed in old galaxies, this mild evolution hints for a growth in the stellar mass of the red sequence, either coming from the blue-peak galaxies in which star formation is truncated by some physical process, or by \"dry mergers\" of smaller red, gas-poor galaxies. However, according to \\citet{Bundy05,Bundy06}, dry mergers cannot be the leading mechanism in the history of galaxy assembly because of the weak dependence on the environment, in contrast to what expected. Indeed the majority of quiescent galaxies seem to be assembled by a mechanisms that depends on their mass rather than the environment, as the merger rate does not seem to increase with environment density \\citep{Bundy06}.  \\textsf{Putting Dynamics and Photometry together.} How can we disentangle the above scenarios? Comparing theoretical predictions to observational data concerning the light and hence mass profiles, velocity, SEDs, magnitudes, colors, line strength indices, and associated gradients. In this context, spectro-photometric models of galaxies have long been the key tool to investigate how galaxies formed and evolved with time.  Consequently, an impressive number of chemo-spectro-photometric models for ETGs have been proposed. To mention a few among the recent ones, we recall \\citet{Bressan94},\\citet{Vazdekis97},\\citet{Tantalo96a},\\citet{Gibson97}, \\citet{Kodama97},\\citet{Fioc97b},\\citet{Tantalo98b},\\citet{Pipino04}, and \\citet{Piovan06,Piovan06b}. Nearly all these models simulate a galaxy and its evolution adopting the point source approximation, in which no morphological structure and no dynamics are considered.  In parallel to this line of work, many fluid-dynamical models in n-dimensions (from 1 to 3) and with multi-phase descriptions of the gaseous component were developed for galaxies of different morphological type. To mention a few, we recall \\citet{Theis92},\\citet{Ferrini93},\\citet{Samland97},\\citet{Boissier99a,Boissier99b}, \\citet{Samland00},\\citet{Samland01},\\citet{Berczik03}, and \\citet{Immeli04}. For some of them the spectro-photometric aspect of the models was also investigated with successful results.  Finally, there are the  NB-TSPH simulations with dark and baryonic matter, the hydrodynamic treatment of the baryonic component, and even multi-phase descriptions of the gas. The NB-TSPH simulations are one of the best tools to infer the 3-D structure of ETGs, to follow the temporal evolution of the dynamical structure, the stellar content, and the chemical elements. Recent models of this type at different level of complexity are by \\citet{Chiocar02},\\citet{Kobayashi04a,Kobayashi04b,Kobayashi05}, \\citet{Merlin06,Merlin07}, and \\citet{Scannapieco06a,Scannapieco06b}. So far the corresponding photometric properties of the models are either left aside or treated in a very rudimentary way.  It follows from these considerations that the ideal tool to develop would be the one folding together NB-TSPH simulations and chemo-spectro-photometry to generate 3-D chemo-dynamical, spectro-photometric models of galaxies. This would allow us to simultaneously predict and discuss both the structural properties related to dynamical formation process and the spectro-photometric ones related to the stellar content in a self consistent fashion, and hopefully to cast light on the above issues. Therefore, we have taken the cosmo-chemo-dynamical models of galaxies calculated with the Padova Code [G\\textsc{al}D\\textsc{yn}] (see details below) from which we get the star formation (SFH), the chemical enrichment (Z(t)) histories, and the structure of the simulated galaxy. The output of these models is fed into the Padova photometric code [S\\textsc{pe}C\\textsc{ody}], which generates the spectral energy distribution (SED) of the whole galaxy. From this SED we derive the absolute magnitudes, colors, indices, etc., in a chosen photometric system. This allows us to determine the rest-frame and cosmological evolution of magnitudes and colors for the set of models at our disposal.  \\textsf{Aims and plan of the paper.} The purpose of this study is to validate the whole procedure before applying it to a set of simulations under preparation and/or to extensive study of observational data. The outline of the paper is as follows. Section~\\ref{dynsim} describes the dynamical NB-TSPH simulations of ETGs. Section~\\ref{specody} describe the photometric package used to get the rest-frame magnitudes and colors of the galaxy models in a photometric system (some details are given for two of those, namely the Bessell-Brett and the Sloan Digital Sky Survey (SDSS). We also we present a study of the color-magnitude diagram (CMD) of the stellar populations of the model galaxies. Section~\\ref{cosmev} presents a multi-wavelength study of the optical and near-IR high-$z$ photometric properties of the ETGs and compares the results with a sample of galaxies selected from COSMOS and GOODS surveys. In Section~\\ref{chapim} we describe the  method followed to derive 2-D artificial images, starting from the 3-D model galaxies. These images resemble observational data and can be analyzed in a similar manner. Isophotal analysis with aid of the Fourier and S\\'ersic technique is applied to derive some structural properties of the model galaxies. We obtain the morphological and structural parameters and compare them with the data for a sample of elliptical galaxies selected from  SDSS. This allows us to establish the consistency of the models with photometric data. In Section~\\ref{scalrel} we use the parameters derived from the surface photometry to investigate  the Kormendy scaling relation. Finally, in Section~\\ref{concl}, we discuss some unsettled issues that require future work and present some general, conclusive remarks. ",
        "conclusions": "\\label{concl}  We have presented a package of numerical codes to compute the spectroscopic and photometric properties of model galaxies by combining the evolutionary population synthesis technique with the 3-D geometrical structure of the galaxy and its history of star formation and chemical enrichment, provided by NB-TSPH simulations. The tool is very flexible in the way input libraries of evolutionary tracks, isochrones, SSPs and important physical laws such as the initial mass function, star formation rate, and metallicity enrichment can be changed, tested, and added to the database for future use. Finally, it can be adapted to any photometric system currently in use. The method has been tested so far on three models of ETGs evolved within different cosmological backgrounds and the analysis has been done in the following photometric systems: Bessell-Brett, SDSS, COSMOS, and GOODS filters among all the ones we have at disposal. Our main results may be summarized as follows:  \\textsf{1.} The application of the  tool to three model galaxies at our disposal, allow us to compute the SED, magnitudes, and colors as function of time and redshift, together with the evolutionary and cosmological corrections. With the aid of it, we analyzed samples of ETGs taken from the COSMOS and GOODS databases, and made a qualitative and quantitative comparison between theoretical results and observational data. For both data-sets, we find that the simulated colors for the different cosmological scenarios follow the general trend over the entire range of redshift considered and are in good agreement with the data up to $z \\sim 1$, above which the number of observed ETGs falls abruptly.  In conclusion, within the redshift range considered, all the simulated colors reproduce quite well the observational data.  \\textsf{2.} We have also generated synthetic 2-D images of the galaxy models in a given photometric system. These synthetic images can be processed with the same algorithms used to analyze real images to derive the structural and morphological parameters, e.g. the galaxy's $R_e$ and the luminosity within this, the shape indices through Fourier and S\\'ersic analysis, the color profiles, and the radial profiles of most of the parameters that define the structure of galaxies. We find that the luminosity profiles of the model galaxies at $z = 0$ can be reasonably fitted with a S\\'ersic $R^{1/n}$ law. The evaluation of the $R_e$ in the photometric bands of SDSS shows the same dependence on the passband wavelength range of the observational data.  Furthermore, the isophotes are well approximated by ellipses, with only a weak radial variation in position angle and ellipticity. Small but significant deviations from perfect ellipses are also measured. In general, we can recover properties that resemble those of observed galaxies.  \\textsf{3.} In addition to that, we looked at the Kormendy relation, one of the Scaling Laws of ETGs, for which all theoretical counterparts of observational data were available. The theoretical luminosities and effective radii of the models are consistent with the archival data from the SDSS for a sample of ETGs.  \\noindent From the above results we can conclude that the package provides good results that permit to reasonable recover the observational properties of ETGs. Owing to the small number of galaxy models to our disposal, neither statistical generalization of the results is possible, nor the underlying cosmological background can be tested. We have to wait for a more complete library of model galaxies in different cosmological scenarios to better elucidate the evolution of photometric properties and scaling laws, and to provide insight into the meaning of the correlations (or lack thereof) between shape, kinematics, and photometry of ETGs."
    },
    "1005/1005.2177_arXiv.txt": {
        "abstract": "{For the evolution of large-scale structures in the Universe, it is unclear whether active galaxies represent a phase that each galaxy undergoes, and whether and to which extent the evolution of black holes at their centers is important. Binary black hole (BBH) systems may play a key role in our understanding of the above questions.} {We investigate the Caltech-Jodrell Bank flat-spectrum (CJF) sample to identify evidence that supports the merger-driven evolution scheme of active galaxies, and search for tracer-systems of active galactic nuclei (AGN) evolution and possible BBH candidates. We discuss the strength of and uncertainty in the evidence and formulate a set of selection criteria to detect such tracer-systems.} {We conduct an extensive literature search for all available  multiwavelength data, particularly in the optical and infrared regime, and morphological information about the CJF sources. We perform a statistical analysis of the properties of this sample.} {We find 1 ULIRG (Mrk 231) included in the CJF prototype of a transitory system. In total, 28.6\\% of the CJF sources with $z\\le0.4$ are distorted or have a companion. Given the unbiased sample used here, this provides strong evidence of the ubiquity of the merger phenomenon in the context of active galaxies. We find a correlation between the radio and the near-infrared luminosity of the high-luminosity sources, interpreted in the context of the interplay between a star-formation and AGN component. We find a connection between variability and evolutionary transitory systems, as selected on the basis of their near-infrared colors. We select 28 sources that trace the different evolution phases of an AGN, as well as a number of the most promising BBH candidates. We find 4 sources with almost periodical variability in the optical and radio on similar timescales.} {} ",
        "introduction": "\\label{sec:intro}  As one of the most prominent predictions of the theory of General Relativity, black holes have been at the center of astronomical research for many years. There is a large amount of indirect evidence that most, if not all, galaxies host a supermassive black hole (SMBH) in their center (e.g., \\citealt{Blandford1986}).  It is widely believed that this SMBH is the main constituent of activity in active galactic nuclei (AGN) (for a review see e.g., \\citealt{Begelman1984}). Although the mechanism that triggers this activity remains unknown, it has been argued by many authors (e.g., \\citealt{Hutchings1983}; \\citealt{Sanders1988}; \\citealt{Hopkins2006}) that galactic mergers should play a predominant role in these processes.  In the context of the hierarchical model (e.g., \\citealt{Frenk1988}), galaxies with a binary black hole (BBH) at their center should be common systems (e.g., \\citealt{Merritt2002}; \\citealt{Haehnelt2002}).  These systems, although predicted in theory (\\citealt{Begelman1980}; for a review on the evolution of BBH see \\citealt{Merritt2005}) have been difficult to study. Only a few of them have been directly observed (e.g., \\citealt{Hudson2006}; \\citealt{Komossa2003}; \\citealt{Rodriguez2006}), while in a number of galaxies the presence of a BBH has been discussed (for an overview, see \\citealt{Komossa2006}). Active galaxies hosting BBH will be primary candidates for the future gravitational wave space interferometer L.I.S.A. (e.g., \\citealt{Sesana2005}).  It is important to identify additional BBH candidates. In this paper, we use the Caltech Jodrell-Bank Flat-spectrum (CJF) flux-limited sample of 293 radio-loud active galaxies to search for these systems, while looking for evidence of the merger-driven evolution scheme for AGN and identifying AGN at different evolutionary stages. This is achieved by means of a multiwavelength analysis of the sample.  This paper is organized as follows. In Sect. \\ref{sec:cjf}, we describe the CJF sample. In Sect. \\ref{sec:bbh}, we describe the different states of BBH systems, their properties, and the criteria to search for them in a large sample of sources. In Sect. \\ref{sec:criteria}, we present the criteria, tracers, with which we investigate the relevance of mergers in AGN. In Sect. \\ref{sec:data}, we present multiwavelength data for the CJF sample. In Sect. \\ref{sec:results}, we present the analysis of the data presented in Sect. \\ref{sec:data}. In Sect. \\ref{sec:tracing}, we present the candidates for different evolutionary stages and BBH systems. We conclude with a discussion in Sect. \\ref{sec:discussion}, a summary, and an outlook. Throughout the paper, we assume the cosmological parameters $H_{0}=71$ $km s^{-1} Mpc ^{-1}$, $\\Omega_{M}=0.27$, and $\\Omega_{\\Lambda}=0.73$. ",
        "conclusions": "\\label{sec:conclusions} The focus of this work has been twofold: (i) to investigate the evolutionary connection between active galaxies and merger events, and (ii) identify AGN candidates at different evolutionary stages, including binary black hole system candidates. To this end, we collect and present multiwavelength information about the CJF sources.  From the previous analysis, it becomes apparent that mergers or merger-induced effects (starburst activity) are ubiquitous in active galaxies. With almost 30\\% of our sample showing evidence of interactions at z$<0.4$ and with indications that more such systems exist, although not yet observed, mergers appear to play a predominant role in the evolution not only of galaxies but AGN as well. We note here the importance of the absence of any biases towards transitory systems in the selection criteria of our sample. This supports a direct link between merger events and active galaxies in general. Furthermore, we show that multiwavelength variability may be linked to transitory systems. Assuming that multiwavelength variability originates in binary black hole systems, this provides an additional link between merger events, binary black holes, and active galaxies, as well as a possible identifier of these sources.  Following the merger-driven evolution scheme, we have also presented a first subsample of 28 CJF sources that are good tracers of evolution in AGN. We find at least 14 promising BBH candidates at different points in their evolution (5 in an unbound and 9 in a hard state). Classification of these sources cannot be unambiguous, given the lack of complete information. Additional observations are needed to clarify their true nature.  Following arguments from both analytical and numerical simulations, it is anticipated that binary black holes should form during merger processes. We show that at least 13.6\\% of the CJF sources exhibit variability at different wavelengths (not including IDV), most of them being almost periodic. We show that most of the previous BBH candidates are highly variable across the spectrum, showing the tight connection between the two phenomena. The morphology and evolution properties of the AGN jets can also be linked to the evolution scheme and BBH systems. The morphology of CJF jets and its evolution will be studied in detail in a future paper.  More specifically, we find that: \\begin{itemize} \\item one CJF source (Mrk 231) has been classified as an ULIRG. Its properties fit well in the merging paradigm of AGN. We also note that the similar source 1146+596, which has a high luminosity in the near-infrared, is probably a merging system. Both sources are primary BBH candidates. \\item a correlation between radio luminosity and NIR luminosity exists. This is qualitatively explained in terms of reradiation of the non-thermal AGN emission. Low luminosity AGN appear to have a near-constant NIR luminosity. \\item 4.4\\% of the sample exhibit distortions in their optical images, signs of ongoing or recent interaction, and that 3.1\\% possibly have a companion. For $z<0.4$, we have found that 28.6\\% of the CJF sources have either distortions or a companion, in excellent agreement with \\citet{Hutchings1983} and \\citet{Surace1998}. \\item the above listed conclusion, combined with the unbiased selection criteria of the CJF sample towards transitory systems, imply that merger events are indeed important to the evolution of active galaxies \\item multiwavelength variability appears to be linked to systems undergoing a transition from a thermal to non-thermal emission-dominated spectrum. This may be connected to the presence of a bound binary black hole in the nucleus of these systems. \\item following the categorization of \\citet{Lotz2008}, we have selected the best candidate sources from our sample for different evolutionary phases. We have also selected sources as the best candidates in our sample for BBH systems. \\end{itemize}  More in-depth optical and infrared observations are needed, to unambiguously distinguish between different phenomena. High quality infrared data (mid and far) of our objects would provide us with the means to pinpoint the true position of each object on the assumed evolution track of galaxies. Alternatively, high resolution observations of those sources with multiwavelength variability would help us to model the possible existence of a binary black hole in these systems."
    },
    "1005/1005.0344_arXiv.txt": {
        "abstract": "We summarize observations with the $Spitzer$ Infrared Spectrograph (IRS) of 571 starbursts (strong PAH emission features), 128 obscured AGN (strong silicate absorption), and 39 unobscured AGN (silicate emission).  Sources range in luminosity from 10$^{8}$ to 10$^{14}$ \\ldot~ and continuously in redshift for 0 $<$ z $<$ 3.  The most luminous starbursts and AGN evolve as (1+z)$^{2.5}$ to z $\\sim$ 2.5; no clear evidence is found that this evolution ceases beyond z = 2.5.  Dust obscuration in starbursts is determined by comparing PAH luminosity with ultraviolet luminosity and indicates severe obscuration in most starbursts, even those selected in the ultraviolet; the median ratio (intrinsic ultraviolet/observed ultraviolet) is $\\sim$ 50 for infrared selected starbursts and $\\sim$ 8 for ultraviolet selected starbursts.  Obscuration increases with bolometric luminosity, but starbursts which appear most luminous in the ultraviolet are those with the least obscuration.  This result indicates that extinction corrections are significantly underestimated for ultraviolet selected sources, suggesting that galaxies at z $>$ 2 are more luminous than deduced only from rest frame ultraviolet observations. ",
        "introduction": "At the same time that the $Spitzer$ mission was being conceived as NASA's final Great Observatory (originally as SIRTF), results from the Infrared Astronomical Satellite (IRAS) demonstrated that the most luminous galaxies in the universe are dusty and heavily obscured, with most of their luminosity arising from optical and ultraviolet luminosity which is reradiated by the absorbing dust \\citep*{soi87}.  A fundamental motive of $Spitzer$ was to track the evolution of these luminous galaxies through the history of the Universe.  In what follows, we review the success of $Spitzer$ in achieving this goal.   Our conceptual sketch of such an ultraluminous infrared galaxy (ULIRG) is shown in Figure 1 and described in the caption.  Undoubtedly, starbursts are present to some extent in all such sources.  The starburst ULIRGs are the most fertile star forming regions in the Universe.  However, the observed characteristics are determined both by the luminosity of the starburst and the luminosity of dust heated by a central active galactic nucleus (AGN).  Here, we separate the two phenomena to track separately the evolution of starbursts and AGN.  The plots and results that we show below derive only from IRS low resolution observations.  More detailed diagnostics can be made for those sources with high resolution data \\citep[e.g.][]{far07,ber09}, but the heavy extinction in the most obscured sources can hide even infrared emission lines arising near the central AGN.  For example, the presence of only blueward asymmetries on infrared forbidden lines from heavily obscured ULIIRG AGN indicates that the obscuring dust is thick, even at the scale of the narrow line region of AGN \\citep{spo09}. \\begin{figure}[!ht] \\figurenum{1} \\includegraphics[scale=0.25]{color.eps} \\caption{Sketch of the assumed model for a luminous, dusty galaxy.  Starbursts can be found within the large torus containing molecular clouds and dust; each starburst immersed in a molecular cloud is surrounded by a photodissociation region (PDR) which produces strong PAH emission features.  Any active galactic nucleus (AGN) heats dust close to the nucleus.  (Note logarithmic size scale at bottom.) The starbursts will appear the same from any viewing direction, but the appearance of the AGN depends on whether it is viewed face on or through the torus.  If face on, the hottest dust is seen directly and silicate emission is seen.  If seen through the torus, the dust continuum from the hot dust is absorbed within cooler dust and produces silicate absorption.  The final observed spectrum is some combination of PAH, silicate absorption, or silicate emission depending on the relative luminosities of starbursts and AGN, and on the viewing direction for the AGN. } \\end{figure} In general, the new IRS results confirm the initial demonstration by \\citet{gen98} that strong emission from polyclyclic aromatic hydrocarbons (PAH) is associated with starbursts, but not with AGN.  The photodissociation region at the interface between the surrounding molecular cloud and the HII region of a starburst provides the appropriate physical conditions for strong PAH excitation and emission \\citep{pee04}.  The hundreds of low resolution spectra accumulated by the IRS give an empirical result that \"pure\" starbursts - sources with no evidence in any spectral region of an AGN - always have equivalent width for the 6.2 \\um PAH feature, EW(6.2 \\ums), to be EW(6.2 \\um)$>$ 0.4 \\um \\citep*{bra06,hou07,wee09a,sar10}.  Conversely to starbursts, there are large numbers of sources with IRS spectra which are \"pure\" AGN, having no evidence in any spectral region of an observationally significant starburst; these have EW(6.2 \\ums)$<$ 0.1 \\ums.  We have used these simple observational criteria to derive the samples of starbursts and AGN presented below, as explained in \\citet{hou07} and \\citet{wee08,wee09a,wee09b}. \\begin{figure}[!ht]  \\figurenum{2} \\includegraphics[scale=0.95]{Weedman_D_2.eps} \\caption{Representative rest frame spectra observed for the most luminous dusty galaxies at z = 2.5.  The top spectrum represents the brightest type 1 AGN and is featureless at these rest wavelengths, but silicate emission begins to increase at the longest wavelengths toward a peak at $\\sim$ 10 \\ums.  Such sources are optically bright and have known redshifts, so IRS spectral measures of the dust continuum are made at 7.9 \\ums.  The center spectrum is a heavily obscured AGN with an apparent spectral peak at 7.9 \\um arising because of deep silicate absorption at longer wavelengths and other absorption features at shorter wavelengths.  The dust continuum is measured at this 7.9 \\um peak.  The lowest spectrum is a starburst with strong PAH features.  Starburst luminosity is measured by the peak of the strong 7.7 \\um PAH feature.  (The illustrated spectra have higher S/N than individual sources actually observed at such redshifts; spectra which are shown derive from brighter, lower redshift examples.)} \\end{figure} Our primary motivation for using the simple observational measurement of peak flux densities for the PAH 7.7 \\um feature in spectra of \"pure\" starbursts or the 7.9 \\um spectral peak for \"pure\" obscured AGN is to have reliable and self-consistent data for the large numbers of sources discovered by $Spitzer$ at z $\\ga$ 2, for which S/N is poor, and deconvolution of the total flux within spectral features is uncertain.  To measure total infrared luminosities, we have used as an initial calibration a limited number of ULIRGs or starbursts having IRAS fluxes and IRS spectra, so that L$_{bol}$ can be estimated as in \\citet{san96}. These calibrations, from \\citet{wee09b}, are given in the captions of Figures 3 and 4.  This estimate of L$_{bol}$ arises only from the IRAS wavelengths, with an assumed correction for colder dust components, so our calibration will be improved as more far infrared and submillimeter observations accumulate.  We use this L$_{bol}$ to calibrate the star formation rate (SFR) according to the precepts of \\citet{ken98}, as given in the caption to Figure 3. \\begin{figure}[!ht]  \\figurenum{3} \\includegraphics[scale=0.95]{Weedman_D_3.eps} \\caption{ Starbursts with IRS measures of PAH luminosity, defined as $\\nu$L$_{\\nu}$(7.7$\\mu$m) for the peak f$_{\\nu}$(7.7$\\mu$m), in units of ergs s$^{-1}$.  Line is the formal fit to evolution of the most luminous sources over all redshifts, which has form (1+z)$^{2.5}$.  Using our adopted calibration, log L$_{bol}$ = log $\\nu$L$_{\\nu}$(7.7$\\mu$m) - 32.80 for L$_{bol}$ in \\ldot~and $\\nu$L$_{\\nu}$(7.7$\\mu$m) in ergs s$^{-1}$;  this yields that log(SFR) = log$\\nu$L$_{\\nu}$(7.7$\\mu$m) - 42.57 for SFR in \\mdot.} \\end{figure} All sources shown below were observed with the IRS, described in \\citet{hou04}, using only the low resolution modules, giving spectral coverage from $\\sim$5\\,\\um to $\\sim$35\\,\\ums.  For all spectra which we analyzed, spectral extractions were done with the SMART analysis package \\citep{hig04}, usually with the improved \"optimal extraction\" procedure in \\citet{leb10}\\footnote{http://isc.astro.cornell.edu/IRS/SmartRelease}. ",
        "conclusions": ""
    },
    "1005/1005.3084_arXiv.txt": {
        "abstract": "Correlations between stellar properties and the occurrence rate of exoplanets can be used to inform the target selection of future planet search efforts and provide valuable clues about the planet formation process. We analyze a sample of 1266 stars drawn from the California Planet Survey targets to determine the empirical functional form describing the likelihood of a star harboring a giant planet as a function of its mass and metallicity. Our stellar sample ranges from M dwarfs with masses as low as 0.2~\\msun\\ to intermediate-mass subgiants with masses as high as 1.9~\\msun. In agreement with previous studies, our sample exhibits a planet-metallicity correlation at all stellar masses; the fraction of stars that harbor giant planets scales as $f \\propto 10^{\\bbest \\rm [Fe/H]}$. We can rule out a flat metallicity relationship among our evolved stars (at 98\\% confidence), which argues that the high metallicities of stars with planets is not likely due to convective envelope ``pollution.'' Our data also rule out a constant planet occurrence rate for [Fe/H]~$< 0$, indicating that giant planets continue to become rarer at sub-Solar metallicities. We also find that planet occurrence increases with stellar mass ($f \\propto M_\\star$), characterized by a rise from 3\\% around M dwarfs (0.5~\\msun) to 14\\% around A stars (2~\\msun), at Solar metallicity. We argue that the correlation between stellar properties and giant planet occurrence is strong supporting evidence of the core accretion model of planet formation. ",
        "introduction": "Mass and chemical composition are key quantities in the formation, evolution and fate of stars. A star of a given age is, to first order, characterized by these two physical parameters, and the influences of mass and metallicity extend to the formation and evolution of planets \\citep{johnson09rev}. Even the first handful of exoplanet discoveries revealed that the likelihood of a star harboring a planet was closely tied to stellar iron content, or metallicity \\feh\\ \\citep{gonzalez97}. Subsequent studies of larger samples of stars using uniform spectroscopic modeling techniques found that giant planet occurrence increases sharply for stellar metallicity in excess of the Solar value, rising from 3\\% for [Fe/H]~$\\lesssim 0$ to 25\\% for [Fe/H]~$> +0.4$ \\citep[][hereafter FV05]{santos04,fischer05b}.  In addition to informing models of planet formation \\citep{ida05a,mordasini09,johansen09}, the planet-metallicity correlation (PMC) has provided a guide for the target selection of subsequent planet searches. The {\\it Next 2000 Stars} (N2K) and Metallicity-Biased CORALIE surveys leveraged the higher metallicities of their samples to detect large numbers of close-in planets, many of which transit their host stars and thereby yield key insights into the interior structures of Jovian exoplanets \\citep{fischer05a,bouchy05,johnson06,moutou06}. Indeed, studies of known transiting planets have revealed evidence of a correlation between planetary core mass and the metallicity of their host stars \\citep{sato05,torres08,guillot06,burrows07c}.  While the first planet detections yielded a definitive correlation between giant planet occurrence and stellar metallicity, until recently very little was known about the effects of stellar mass \\citep{laws03}. The first Doppler-based planet surveys concentrated primarily on stars with masses similar to the Sun, both because it was desirable to find Solar System analogs and because Sun-like stars make excellent planet-search targets. Compared to more massive stars, dwarfs with masses within  $1.0\\pm0.2$~\\msun\\ are relatively numerous, have cool atmospheres and slow rotational velocities (\\vsini~$ \\lesssim 5$~\\ks). The latter two features result in a high density of narrow absorption lines in the spectra of Sun-like stars, which is ideal measuring for stellar Doppler shifts to high precision.  Stars at the lower end of the mass scale (the K and M stars) are even more numerous than the Sun and they also display large number of narrow absorption features in their spectra. However, most low-mass stars are optically faint ($V_{mag} \\gtrsim 10$) and are thus not included in large numbers in most Doppler surveys. The faintness of late-K and M-type dwarfs can be overcome by using larger telescopes \\citep{butler04,bonfils05b}, and more recently by observing at infrared wavelengths \\citep{bean10}. Despite the small numbers of M dwarfs thus far monitored by Doppler surveys, one result has become apparent: M dwarfs harbor Jovian planets very infrequently. Only eight systems containing one or more giant planets have been found among the $\\approx 300$ M dwarfs on various Doppler programs \\citep{johnson10a, nader10}.  It was originally thought that the paucity of Jupiter-mass planets around M  dwarfs was due to a metallicity bias among nearby, low-mass stars \\citep{bonfils05a}. However, a recent study by \\citet{johnson09b} revealed that M dwarfs likely have the same metallicity distribution as Sun-like stars, and stars with masses $M_\\star < 0.5$~\\msun\\ are 2-4 times less likely than Sun-like stars to have a Jupiter \\citep{johnson07b, johnson10a}.  At the other end of the mass scale, the problems inherent to massive, early-type stars can be overcome by observing targets at a later stage of their evolution \\citep{hatzes03, setiawan05, sato05, reffert06, johnson07,nied07, liu08, dollinger09}. Once stars exhaust their core hydrogen fuel sources they move off of the main sequence, become cooler, and shed a large fraction of their primordial angular momentum \\citep{gray85, donascimento00}. The effects of stellar evolution transform a 2~\\msun\\ star from an A-type dwarf with \\vsini~$\\sim 100$~\\ks\\ and \\teff~$ = 8200$~K, to a K-type subgiant or giant with \\vsini~$< 2$~\\ks\\ and \\teff~$\\approx 4800$~K \\citep{demedeiros97, girardi02,sandage03}. Surveys of ``retired'' massive stars have resulted in the discovery of $\\approx 30$ Jupiter-mass planets with well-characterized orbits \\citep[see e.g. Table 1 of][]{bowler10}.  Using a sample of stars spanning a wide range of masses, \\citet{johnson07b} measured a positive correlation between stellar mass and the fraction of stars with detectable planets\\footnote{In that study, ``detectable planets'' were defined as having \\msini~$ > 0.8$~\\mjup\\ and $a < 2.5$~AU.}. In a related study, \\citet{bowler10} measured a planet occurrence rate of $26^{+9}_{-8}$\\% among a uniform sample of 31 massive subgiants. Furthermore, based in part on a study of planets around K giants in nearby open clusters, \\citet{lovis07} found that the average planet mass increases as a function of stellar mass, indicating that gas giant planets become either more massive on average, or more numerous (or both) with increasing stellar mass \\citep[see also][]{bowler10}.  The observed correlation between stellar mass and the occurrence of detectable planets, like the PMC before it, has added an important new variable to models of planet formation. While the Sun and Solar-mass stars serve as important benchmarks for understanding the formation of our own planetary system, successful, generalized planet formation theories must now account for the effects of stellar mass, and presumably by extension, disk mass \\citep{laughlin04, ida05a, kennedy08, currie09}.  While previous studies have uncovered the existence of a positive correlation between stellar mass and planet occurrence, it is important to understand the underlying functional form of the relationship. For example, it would be advantageous to know whether the correlation is purely linear, or if it can instead be better described as some other functional form. Besides informing theories of planet formation, an improved understanding of the relationship between planet occurrence and stellar mass will also help guide the target selection of future surveys, and aid in the interpretation of results of current and future planet-search efforts. Just as some previous Doppler surveys biased their target selection toward high-metallicity stars to increase their yield, future direct-imaging, astrometric and Doppler surveys may benefit from concentrating on more massive stars.  This strategy has paid off for one high-contrast imaging survey, resulting in the detection of three giant planets around the A5 dwarf HR\\,8799 \\citep{marois08}. Another example of an imaged planet is Fomalhaut\\,b, which is a giant planet ($\\lesssim3.3$~\\mjup) orbiting just inside of a debris disk of an A3V star \\citep[][]{kalas08,chiang09}. Even in the cases when surveys do not yield detections, proper interpretation of null results requires knowledge of the expected number of detections \\citep[e.g.][]{nielsen09}.  Ascertaining the underlying form of the dependence of giant planet occurrence on stellar mass requires a larger sample than used in previous studies. Since the publication of \\citet{johnson07b} a sample of 240 new intermediate-mass subgiants have been added to the California Planet Survey (CPS) at Keck Observatory \\citet{johnson10b}. At the low-mass end, two new giant planets have been  discovered among the CPS Keck sample of M dwarfs \\citep{johnson10a}. Improvements in our ability to estimate the metallicities of M dwarfs and massive evolved stars have provided vital information about how to properly isolate the effects of stellar mass from the known effects of stellar metallicity. With these tools at hand, we are now poised to make an updated evaluation of the relationship between stellar mass and planet occurrence.  Our paper is organized as follows. In \\S~\\ref{sample} we present the characteristics of our three primary samples, including low-mass M dwarfs from Keck observatory; ``Sun-like'' late-F, G and K (FGK) dwarfs from the main CPS sample; and massive, evolved stars from the Lick and Keck subgiant surveys. In \\S~\\ref{metallicity} we examine the separate effects of mass and metallicity on planet occurrence. In \\S~\\ref{fraction} we present our Bayesian inference technique of measuring correlations between planet occurrence and stellar characteristics and we provide the best-fitting parameters for the measured relationship in \\S~\\ref{results}. We compare our results with previous work in \\S~\\ref{comparison}. Finally, we summarize our key results and discuss our findings in the context of the current theoretical understanding of planet formation in \\S~\\ref{summary}.  \\begin{figure} \\epsscale{1.2} \\plotone{f1.eps} \\caption{H-R diagram showing our stellar sample (filled circles) and planet host stars (open circles). The solid line is the polynomial relationship describing the mean \\emph{Hipparcos} main sequence \\citet{wright04}. The excess scatter seen about the lower main sequence $B-V > 1.2$ is primarily due to errors in the published  V-band magnitudes and/or parallaxes for those faint stars. \\label{fig:hr}} \\end{figure} ",
        "conclusions": "\\label{summary}  We have used a large sample of planet-search target stars and planet detections from the CPS to study the correlation between stellar properties and the occurrence of giant planets ($K > 20$~\\ms) with $a \\lesssim$2.5~AU. We have derived an empirical relationship describing giant planet occurrence as a function of stellar mass and metallicity, given by  \\begin{eqnarray} f(M_\\star,{\\rm [Fe/H]}) &=& \\cbest \\pm 0.01 \\times (M_\\star/M_\\odot)^{\\abest \\pm 0.3} \\nonumber \\\\ &\\times& 10^{\\bbest \\pm 0.2 \\rm [Fe/H]}. \\end{eqnarray}  Our understanding of planet formation is presently dominated by two theories: core accretion \\citep[e.g.][]{pollack96} and disk instability \\citep{boss97}. The core accretion model is a bottom-up process, by which protoplanetary cores are built up by the collisions of smaller planetesimals. Once the core reaches a critical mass of roughly 10~$M_\\oplus$, it rapidly accretes gas from the surrounding disk material. The disk instability mechanism is a top-down process whereby giant planets form from the gravitational collapse of an unstable portion of the protoplanetary disk.  Both models depend on the existence of a massive gas disk, a portion of which forms the bulk of the final mass of the Jupiter-like planet. However, since the inner gas disks of protoplanetary disks disperse on timescales of 3-5~Myr, the process of planet formation is a race against time \\citep[e.g.][]{pickett04}. In this race the disk instability model holds a major advantage over the core accretion model because, under the right conditions, planets can form from disk collapse in a mere thousands of years, compared to of order Myr timescales required for core accretion.  The disk instability model predicts that there should be no dependence on planet formation and physical stellar properties. The simulations of \\citet{boss06} showed that giant planets should readily form in even low-mass protoplanetary disks, and that in general giant planets should form efficiently via disk instability independent of stellar mass. The disk instability model also predicts that planet formation should also be independent of disk metallicity \\citep{boss02}. Indeed, \\citet{cai06} and \\citet{meru10} show that the efficiency of disk instability to form giant planets {\\it decreases} with increasing metallicity.  In contrast to these predictions of the disk instability model, we observe a strong dependence between planet occurrence and the physical properties of the star. Assuming that the present-day mass and metallicity of a star reflects the conditions in its protoplanetary disk, then our results suggest that disk instability is not the primary formation mechanism for the giant planets detected by Doppler surveys. Indeed, it has long been recognized in that there are theoretical complications in forming close-in planets via disk instability \\citet[][e.g.]{boley09}. However, the mechanism may be responsible for planets in wide orbits \\citep{dr09}, however see \\citet{kratter10} for complications with this scenario.  An alternative explanation for the observed planet-metallicity correlation is the so-called ``pollution'' model. In this scenario, planet formation actually occurs around stars of all metallicities, and the accretion of gas-depleted protoplanets onto the thin convective layers of stars gives rise to an enhanced stellar metallicity that is actually only \"skin deep\" \\citep{murray02}. \\citet{pasquini07} interpreted the flat metallicity distribution of K giants with planets as evidence for such an effect, since the deepening convective envelopes of evolved stars should dilute any metallicity enhancement of its outer layers. However, we find that the planet fraction among our massive subgiants is described well by a model with a monotonic rise as a function of both mass and metallicity (Section~\\ref{comparison}).  While much attention has been given to evolved, intermediate-mass stars in the investigation of the pollution paradigm \\citep[][FV05]{laughlin97}, stars at the other end of the stellar mass scale provide another proving ground. M dwarfs have deep convective envelopes over their entire lifetimes, and stars with masses below 0.4~\\msun\\ are expected to be completely convective, at least in the absence of strong magnetic activity \\citep[e.g.][]{mullan01}.  The evidence for a PMC among the subgiants, together with a strong PMC seen among the M dwarfs, are highly suggestive that the present-day metallicities of stars are representative of the compositions of their disks during the planet-formation era. One compelling explanation for both the observed PMC and stellar-mass correlation is that the surface density of solids is a key factor in the planet formation process \\citep{laughlin04, ida04, robinson06}. If so, both higher stellar (disk) metallicity and higher stellar (disk) mass can generate the requisite surface density for planet formation.  The relationship between planet formation efficiency and stellar mass/metallicity has been previously studied in the context the  core accretion paradigm \\citep{ikoma00,kornet05,kornet06, kennedy08, thommes08, kretke09,johansen09,mordasini09,dr09}. \\citet{kennedy08} modeled the evolution of the temperature profile at the disk midplanes of stars of various masses, and studied the width of the radial region of disks in which protoplanetary cores form most efficiently. They found that the disks around A-type stars on their descent to the main sequence  have very broad formation regions, and their models predicted a positive correlation between stellar mass and giant planet formation efficiency. Their prediction for the planet fraction is shown as a dashed line in Figure~\\ref{fig:modelfit}, which we have approximated using the polynomial relationship $f(M) = -0.03633 + 0.0138 M_\\star - 0.0060 M_\\star^2$. The agreement between theory and observation is striking.  The interplay between the mass and metallicity of protoplanetary disks is also apparent in the core accretion simulations of \\citet{thommes08}. In their analysis, they simulated disks with a wide variety of masses, viscosities and metallicities. Their models produce gas giants most effectively in disks with a combination of high masses ($M_{\\rm disk} \\gtrsim 0.04$~\\msun) and low viscosities. In their simulations of the effects of metallicity, they found that gas giants can form in Solar-composition disks only if the disk masses exceed $\\approx 0.06$~\\msun, or twice the minimum-mass Solar nebula. This mass threshold decreases to $\\approx 0.03$~\\msun\\ for disks with [Fe/H]~$ = 0.25$. Thus, Thommes et al. showed that there can be a trade-off between the mass and metallicity of a protoplanetary disk in forming giant planets. In the core accretion paradigm, M dwarfs can form giant planets, but only if they have high metallicities. Similarly, even low-metallicity A stars can form massive planets, owing to their more massive disks \\citep[see also][]{dr09}.  Theories of planet formation have progressed in large leaps thanks largely to the rapidly growing sample of exoplanets discovered around other stars. Successful, generalized theories of the origins of planetary systems must account for the observed correlations between planet occurrence and stellar properties. The findings presented herein are the result of more than 15 years of high-precision Doppler monitoring of nearby stars. Additional information will soon pour in from other surveys using techniques such as microlensing \\citep[e.g.][]{dong09,gould10}, astrometry \\citep{boss09}, transits \\citep{kepler,corot1,irwin09} and direct imaging \\citep{sphere,nici,gpi}.  Our results have important implications for the target selection of these future planet surveys. We find that A-type stars harbor planets at an elevated rate compared to less massive stars. However, it is not obvious whether A type stars make the most promising targets for other types of surveys. Where massive stars perhaps hold the most promise is for direct imaging surveys. The first two planetary systems imaged around normal stars were both young A-type dwarfs \\citep{kalas08,marois08}. Are A dwarfs the ideal direct imaging targets? The answer to this question will rely on the results presented herein, along with a careful consideration of the mass, metallicity, luminosity and age distributions of nearby stars; and the orbital and physical properties of planets as a function of stellar mass. This issue will be addressed in a companion paper (Crepp \\& Johnson 2010)."
    },
    "1005/1005.1048_arXiv.txt": {
        "abstract": "{The Herbig Ae/Be star \\object{HD~100546} hosts one of the most extensively studied discs in the southern sky. However, not much is known about this disc's gas and large scale kinematics.} {Our aim is to establish whether the disc is gas-rich, and to study the disc temperature and kinematics.} {We detect and study the molecular gas in the disc at spatial resolution from 7$\\farcs$7 to 18.9$\\farcs$ using the Atacama Pathfinder Experiment telescope\\footnote{This publication is based on data acquired with the Atacama Pathfinder Experiment (APEX). APEX is a collaboration between the Max-Planck-Institut f\\\"ur Radioastronomie, the European Southern Observatory, and the Onsala Space Observatory.}. The lines $^{12}$CO $J=$7--6, $J=$6--5, $J=$3--2, $^{13}$CO $J=$3--2 and [C I] $^3$P$_2$--$^3$P$_1$ are observed, diagnostic of disc temperature, size, chemistry, and kinematics. We use simple parametric disc models and a molecular excitation and radiative transfer code to fit the observed spectral line profiles.} {Our observations are consistent with more than 10$^{-3}$~M$_{\\odot}$ of molecular gas in a disc of $\\approx$400~AU radius in Keplerian rotation around a 2.5~M$_{\\odot}$ star, seen at an inclination of 50$\\degr$. The detected $^{12}$CO lines arise from gas at 30-70~K.  Asymmetry in the $^{12}$CO line emission suggests that one side of the outer disc is colder by 10-20~K than the other, and is not due to pointing offsets, cloud absorption and asymmetry in the disc extent. } {Both low- and high-$J$ $^{12}$CO lines are dominated by the outermost disc regions, and indicate a 400~AU radius. The $^{12}$CO $J=$6--5 line arises from a disc layer higher above disc midplane, and warmer by 15-20~K than the layer emitting the $J=$3--2 line. The existing models of discs around Herbig Ae stars, asuming a B9.5 type model stellar atmosphere overproduce the [CI] $^3$P$_2$--$^3$P$_1$ line intensity from HD~100546 by an order of magnitude.} ",
        "introduction": "\\label{s:intro}  Over the past decade our understanding of the structural and physical properties of discs around young stars has increased from basic theoretical modelling of the spectral energy distributions (SEDs) constrained by observations with no spatial information, to modelling based on not only the SEDs, but also spatially resolved dust observations, like scattered light images and interferometry \\citep{pinte,panic2008,tannirkulam}. Two decades ago, the first submillimetre interferometer observations resolved the molecular gas emission spatially and this allowed major progress in understanding the disc kinematics, structure and chemistry \\citep[e.g.,][]{beckwith1987,koerner1993,dutrey1994}. {A more recent example is a study of one of the brightest discs around a Herbig~Ae star, \\object{HD~163296} shown in \\citet{isella}.} (Sub)millimetre gas and dust emission is the ideal probe of the global disc properties, like size, mass and radial distribution of disc material, because the bulk of the disc mass is located beyond 100~AU from the star, at temperatures of 10-50~K that dominate this part of the spectrum. Disc models which include constraints of both dust and molecular gas observations have stressed the importance of analysing the gas and dust components simultaneously, in the context of a {single disc model with a three-dimensional temperature and density structure} \\citep[][]{wilner,raman,panic2008,panic2009}.  Until recently, observations of rotational transitions of molecules in the submillimetre regime were focused primarily on the low-$J$ emission from $^{12}$CO, up to the $J=$3-2 line \\citep{greaves2000,thi2001,qi2004,thi2004,dent2005}. In two of the brightest and most studied sources, \\object{TW~Hya} and \\object{LkCa~15}, the observations of higher-$J$ transitions of $^{12}$CO, up to J$=$6--5 ($E_{\\rm k}=$116~K), were compared to the low-$J$ lines, providing estimates of the gas temperature in the intermediate-height molecular layer \\citep{vzadelhoff}, crucial ingredients for chemical modelling of discs. These single-dish line spectra were fitted using simplistic disc models, deriving a temperature of 20-40~K in the $^{12}$CO line emitting layers of LkCa~15, and more than 40~K in TW~Hya. \\citet{qi2006} analysed submillimetric interferometer observations of TW~Hya in the context of a disc structure based on an accretion disc model \\citep{calvet2002}. Based on $^{12}$CO $J=$6--5, 3--2 and 2--1 observations, they show that X-ray heating of the gas is efficient in this source, in addition to the stellar radiation field. Such diagnostics of gas heating and ionisation improve our understanding of how the gas content evolves in discs.  The emerging (sub)millimetre facilities in the Southern hemisphere like the Atacama Pathfinder EXperiment (APEX) and the Australia Telescope Compact Array (ATCA) are opening a window towards the star-forming regions of the Southern sky and are well suited to study circumstellar disc emission. These instruments also pave the path for future observations with the Atacama Large Millimetre / Submillimeter Array (ALMA), which will drastically improve our knowledge of disc structure and evolution \\citep[e.g.][]{GD2008}. We use APEX receivers APEX-2a and CHAMP$^+$ to observe the $^{12}$CO $J=$7--6, $J=$6--5, $J=$3--2, $^{13}$CO $J=$3--2 and [C I] $^3$P$_2$--$^3$P$_1$ line emission towards the disc around the young intermediate-mass star HD~100546. A wealth of observations of dust in this bright disc \\citep{waelkens1996,malfait,grady2001,augereau,bouwman2003,acke2006,ardila2007} has motivated us to probe its molecular gas content and kinematics. The chosen transitions are particularly sensitive to the gas in the warm upper layers and kinematics of the outer disc. Our millimetre line observations probe the outer radius and inclination. The existing observational constraints on these parameters in the disc around HD~100546 provide an excellent basis for the analysis of our data. Our observations also provide a bridge toward even higher-$J$ far infrared $^{12}$CO transitions to be observed with the Herschel Space Observatory. HD~100546 is a young B9V type, 2.5~M$_{\\odot}$ star, classified as a Herbig Be star due to its isolation, infrared excess and silicate emission \\citep{the1994,malfait}. With a distance of 103$\\pm$6~pc, measured by Hipparcos, this is one of the nearest Herbig Ae/Be stars. The age of the star is estimated to be greater than 10~Myr \\citep{vdancker}. This makes the presence of circumstellar material intriguing, considering that discs are found to dissipate within 10~Myr in most young stars \\citep[e.g.,][]{hollenbach2000,hernandez2007,hillenbrand2008}. Based on SED modelling, \\citet{bouwman2003} postulate the presence of an inner hole in the disk with a 10~AU radius, possibly caused by a Jupiter-sized planet \\citep[see also][]{acke2006}. Direct evidence of cold disc material at larger radii is provided by ATCA observations of \\citet{wilner} at 89~GHz (3.4~mm) and 2$\\arcsec$ resolution, with the flux of 36$\\pm$3~mJy, values consistent with the 1.3~mm observations of \\citet{henning1998}. They do not detect HCO$^+$ $J=$1--0 line emission and speculate that photodissociation of its progenitor species $^{12}$CO, in the upper disc layers or an overall gas depletion may be the reason for this. Recent spectroastrometric observations of rovibrational $^{12}$CO transitions by \\citet{vdplas} suggest that this molecular species is missing from the inner disc, at least up to 8~AU from the star \\citep[also see][]{brittain2009}.  Scattered light imaging of HD~100546 reveals the disc extending up to 4$\\arcsec$ from the star viewed at an inclination of 50$\\degr$, and an interesting disc structure resembling spiral arms \\citep{pantin,augereau,grady2001}. This structure was interpreted as due to disc perturbation by a companion \\citep{quillen2005} or a warped disc structure \\citep{quillen2006}. Coronographic imaging by \\citet{augereau} shows steep surface brightness profiles in the environment of HD~100546 indicative of optically thin emission in the near-infrared, with surface densities as low as 10$^{-3}$~g~cm$^{-2}$. Their images trace the emission of small dust ($<5\\,\\mu$m), extending out to 800~AU from the star. The authors suggest the presence of an optically thick disc with a 400~AU radius, and an optically thin flattened halo or envelope farther from the star. The scattered light images hint at the presence of gas in the disc that supports the disc vertical structure.  In this work, we detect and study the molecular gas, its kinematics and temperature, in the disc at spatial resolution from 7$\\farcs$7 to 18.9$\\farcs$. In Sect.~\\ref{s:obs&res} we present our observations of $^{12}$CO, $^{13}$CO and [C I] lines. All $^{12}$CO lines are detected, there is a tentative detection of the $^{13}$CO line, while [C I] emission is not detected. We model the spectra in Sect.~\\ref{s:discussion} discussing the implications for the disc size, mass and kinematics. We identify the regions dominating the observed lines, and derive their temperatures. Section~\\ref{s:conclusions} summarises our results. ",
        "conclusions": "\\label{s:conclusions}  We summarise our conclusions as follows:  \\begin{itemize} \\item We present evidence for warm molecular gas associated with the disc around HD~100546, in the regions within 400~AU from the star, successfully separated from more extended material in our CHAMP$^+$ observations; \\item The gas kinematics are consistent with Keplerian rotation around an 2.5~M$_{\\odot}$ star of a disc with a 400~AU radius, viewed at an inclination of 50$\\degr$ from face-on; \\item The $^{12}$CO (6--5)/(3--2) line ratio of 1.1$\\pm$0.6 is higher than measured towards discs around T Tauri stars, likely due to a more efficient gas heating of the disc containing PAHs by the stronger UV radiation from the B9 star; \\item {Our data testify to the significant molecular gas reservoir in the disc, consistent with the total disc masses of more than 10$^{-3}$~M$_{\\odot}$.} We exclude the possibility of a low-density disc and optically thin $^{12}$CO emission. \\item Line asymmetry seen in $^{12}$CO $J=$6--5 and $J=$3--2 lines is explained by a temperature asymmetry, with one side of the disc slightly colder than the other, possibly due to a partial obscuration of one side by a warped inner disc or a high disc rim. We exclude radial asymmetry, midplane density asymmetry and mispointing as possible causes; \\item Our modelling shows that, due to the efficient heating of the disc gas by the star, both low-$J$ and high-$J$ $^{12}$CO lines are dominated by the outermost regions of the disc, though slightly different vertical disc layers with $\\Delta T$=15-20~K; \\item {We find that in `colder' discs where temperatures of the emitting regions are close to 20-30~K in the outer disc -- colder Herbig Ae discs and especially T Tauri discs -- the high-$J$ lines probe a larger extent of the disc, starting from as little as 50~AU}; \\item The non detection of [C I] $J=$2--1 line may indicate the presence of more carbon-ionising photons than assumed in the B9 model atmosphere. \\end{itemize}  Future observations with ALMA will be crucial to characterise the disc around HD~100546, and spatially resolve its kinematics and structure. In particular, these observations will allow a detailed comparison between the spatial distribution of the gas traced by the rotational transitions of $^{12}$CO and its isotopologues, and the dust traced with the millimetre continuum emission. Herschel far-infrared data can probe even higher $J$ $^{12}$CO transitions, as well as [O I] and [C II] lines originating from the disc surface. Being a bright, isolated source suspected to harbour a planet, the disc around HD~100546 is one of the prime targets to probe disc structure in the early planet-building phases."
    },
    "1005/1005.1081_arXiv.txt": {
        "abstract": "The accretion-induced collapse (AIC) of a white dwarf  to form a neutron star can leave behind a rotationally supported disk with mass of up to $\\sim 0.1 \\msun$.  The disk is initially composed of free nucleons but as it accretes and spreads to larger radii, the free nucleons recombine to form helium, releasing sufficient energy to unbind the remaining disk.   Most of the ejected mass fuses to form $^{56}$Ni and other iron group elements.  We present spherically symmetric radiative transfer calculations of the transient powered by the radioactive heating of this ejecta.  We estimate the ejecta composition using nucleosynthesis calculations in the literature and explore the sensitivity of our results to uncertainties in the ejecta kinematics.  For an ejecta mass of $10^{-2} \\msun$ ($3 \\times 10^{-3} \\msun$), the lightcurve peaks after $\\lesssim 1$ day with a peak bolometric luminosity $\\simeq 2 \\times 10^{41} \\ergss$ ($\\simeq 5 \\times 10^{40} \\ergss$); the decay time is $\\simeq 4 \\, (2)$ days.  Overall, the spectra redden with time reaching $U-V  \\simeq 4$ after $\\simeq 1$ day; the optical colors ($B-V$) are, however, somewhat blue.  Near the peak in the lightcurve, the spectra are dominated by Doppler broadened Nickel features, with no distinct spectral lines present.  At $\\sim 3-5$ days, strong Calcium lines are present in the infrared, although the Calcium mass fraction is only $\\sim  10^{-4.5}$.  If rotationally supported disks are a common byproduct of AIC, current and upcoming transient surveys such as the Palomar Transient Factory should detect a few AIC per year for an AIC rate of $\\sim 10^{-2}$ of the Type Ia rate.  We discuss ways of distinguishing AIC from other rapid, faint transients, including .Ia's and the ejecta from binary neutron star mergers. ",
        "introduction": "\\voffset=-1cm \\vspace{0.2 cm} \\label{sec:int} \\label{sec:intro}  As an accreting white dwarf (WD) approaches the Chandrasekhar mass, it suffers one of two fates:  (1) collapse to form a neutron star (NS) or (2) explosion as a Type Ia supernova (SN).  Which of these end states is realized in different astrophysical scenarios is not fully understood (see ~\\citealt{Canal&Gutierrez97} for a review).  The composition of the WD is, however, a critical factor:  the relatively low densities required for electron captures on Ne makes accretion-induced collapse (AIC) of a WD to a neutron star possible, and thus O-Ne WDs, rather than C-O WDs, preferentially undergo AIC (\\citealt{Nomoto84}; \\citealt{Miyaji+80}; \\citealt{Canal+90}; \\citealt{Canal+92}; \\citealt{Gutierrez+96, Gutierrez+05}; \\citealt{Wanajo+09}).  AIC may result both from accretion following Roche-lobe overflow from a non-degenerate companion (the ``single degenerate'' AIC channel; e.g.~\\citealt{Nomoto+79}; \\citealt{Nomoto&Kondo91}), as well following the merger of two WDs in a binary system (e.g.~\\citealt{Yoon+07}).  There is no direct observational evidence for AIC and so its rate remains uncertain.   Numerical simulations of AIC find that it is accompanied by a weak explosion ($\\simeq 3 \\times 10^{49}$ ergs) that synthesizes only a very small amount of Ni ($\\lesssim 10^{-3} \\msun$ in current models) \\citep{WB,Dess06,Ott10}, unless there is a strong, large-scale magnetic field in the progenitor white dwarf \\citep{Dess07,Metzger+08}.  The neutrino-driven wind from the proto-neutron star does produce $\\sim 3 \\times10^{-3} \\msun$ of neutron-rich elements (specifically, matter with an electron fraction $\\lesssim 0.4$; \\citealt{Dess06}).  The low total abundance of such elements in the galaxy implies an AIC rate of $\\lesssim 1 \\%$ of the Type Ia rate, if the calculations of \\citet{Dess06} are representative of AIC in general.   On empirical grounds, \\citet{QianWass07} have suggested that AIC may in fact be the origin of the third peak $r$-process elements.  Because a WD accretes angular momentum as well as mass, it is likely to be rapidly rotating as it approaches $M_{ch}$.  Evidence in favor of this expectation is provided by the recent discovery that the X-ray pulsator RX J0648.0-4418 is an accreting WD with a mass $\\gtrsim 1.2M_{\\sun}$ and a rotational period $P \\sim 13$ seconds \\citep{m09}.  Rapid rotation of the WD prior to collapse would make AIC a strong source of gravitational wave emission \\citep{Ott10}.   An electromagnetic counterpart to such a detection would significantly increase its scientific yield.  More generally, the detection of AIC would provide important new constraints on binary WD evolution and the origin of Type Ia SNe.  New optical transient surveys operating with high cadence and wide fields of view (e.g., Palomar Transient Factory [PTF; \\citealt{Rau+09,Law+09}] and the Pan-STARRs-1 Medium Deep Survey) are sensitive to much fainter transients than previous surveys -- AIC may thus finally be detectable.   In order to increase the probability of such a detection, it is critical to have diagnostics that can isolate AIC among the deluge of transients that are being discovered.  In this paper, we present Monte-Carlo calculations of the optical-infrared lightcurves and spectra of transients produced by AIC.   Our radiative transfer calculations are motivated by the model presented in \\citet{MPQ09-Nickel}.   They argued that the observational signature of AIC can, in some circumstances, be dominated by the ejecta produced by an accretion disk around the newly formed neutron star, rather than by the weak SN that accompanies the AIC itself.  To motivate our radiative transfer calculations, we first summarize \\citet{MPQ09-Nickel}'s calculation of the disk's evolution and nucleosynthesis (\\S \\ref{sec:summary}).  We then describe our radiative transfer calculations (\\S \\ref{sec:fidmodcode}) and the composition and structure of the disk ejecta (\\S \\ref{sec:fidmodics}).   Our main results are presented in \\S \\ref{sec:results}.   We conclude by discussing the key properties of AIC, and comparing them to other predicted faint, rapid transients (\\S \\ref{sec:discussion}).   \\vspace{-0.4cm} \\subsection{Summary of the Model}  \\label{sec:summary}  Depending on the rotation profile of the WD prior to collapse, and in particular the degree of differential rotation (\\citealt{Saio&Nomoto04};\\citealt{Yoon&Langer04};\\citealt{Piro08}), the neutron star formed during AIC is surrounded by a rotationally supported disk \\citep{Dess06, Dess07,Ott10}.  For even modest differential rotation (that does not significantly change the Chandrasekhar mass), the disk mass is $M_D \\sim 10^{-2}-0.1 \\, M_\\odot$.  However, if the WD is uniformly rotating, the disk mass is significantly lower, even if the WD is rotating near break-up \\citep{Ott10}.  Most of the rotationally supported material is initially just outside the proto-neutron star at a radius of $R_D \\sim$ 100 km. The disk is born sufficiently hot that nuclei are dissociated into free nucleons and weak interaction processes have time to change the neutron to proton ratio. Under these conditions, the composition of the disk can be usefully described using the electron fraction \\begin{equation} Y_e = \\frac{n_p}{n_n+n_p} \\end{equation} where $n_p$ and $n_n$ are the proton and neutron densities, respectively. Weak interactions in the disk are dominated by the pair capture reactions \\begin{equation} e^+ + n \\rightarrow \\bar{\\nu}_{e} + p \\label{eq:p} \\end{equation} \\begin{equation} e^- + p \\rightarrow \\nu_{e} + n \\label{eq:n} \\end{equation} where the resulting neutrino emission is also the dominant coolant for the disk material because the photons are trapped and cannot escape on the timescales of interest.  Although the disk is formed very hot, electron degeneracy suppresses the formation of electron/positron pairs.  The resulting electron excess favors electron captures (eq.~[\\ref{eq:n}]) over positron captures (eq.~[\\ref{eq:p}]), thus driving the composition neutron rich (e.g.~\\citealt{Pruet+03}).  As the disk accretes onto the proto-neutron star, the remaining mass spreads to larger radii in order to conserve angular momentum.   As the disk spreads, the temperature drops (for the majority of the mass) and thus the neutrino cooling and weak reactions become less efficient. When the bulk of the disk reaches $R_D \\sim$ 1000 km, the disk has cooled to T $\\sim 10^{10}$K, and weak interactions freeze out \\citep{MPQ09-Neutron}.  Absent a central neutron star -- as is the case for analogous disks formed during binary neutron star  mergers -- weak interactions freeze out when the material is still quite neutron rich, with $Y_e \\sim 0.2-0.4$.   During AIC, however, the proto-neutron star emits copious neutrinos in the first few seconds after formation \\citep{BL}, which irradiate the disk as it spreads.  Because the neutron star is deleptonizing, electron neutrinos dominate the emission and the rate of neutrino captures $\\nu_{e} + n \\rightarrow p + e^{-}$ becomes comparable to the rate of pair captures (eqs.~[\\ref{eq:p}]-[\\ref{eq:n}]).  For a significant fraction of the mass, this increases the electron fraction at freeze-out to $Y_e \\gtrsim$ 0.5.  As the disk continues to spread, the free nucleons in the disk recombine to form heavy nuclei once T $\\sim$ few $\\times 10^9$ K.   The formation of He alone releases $\\sim 7$ MeV per nucleon, which is larger than the gravitational binding energy of the disk when fusion commences ($\\sim 2$ MeV per nucleon).  As a result, the disk is unbound, with its slow viscous spreading transformed into a supersonic outflow that reaches $v \\simeq 0.1$ c.  The end state of the fusion that initiates the disk's destruction is determined by the density, temperature, and electron fraction (the latter because weak interactions have frozen out); the parameters are such that the ejecta is nearly in nuclear statistical equilibrium (NSE). Since there are roughly equal numbers of protons and neutrons, a significant fraction of the unbound disk fuses to form $^{56}$Ni,  the decay of which can power a supernova-like transient once the ejecta expands sufficiently that photons may diffuse out.  In the remainder of this paper, we calculate the lightcurves and spectra of this transient.  \\vspace{-0.5cm} ",
        "conclusions": "Depending on the rotation profile of the WD prior to collapse, and in particular the degree of differential rotation (\\citealt{Saio&Nomoto04};\\citealt{Yoon&Langer04};\\citealt{Piro08}), the neutron star formed during AIC is surrounded by a rotationally supported disk \\citep{Dess06, Dess07,Ott10}.  For even modest differential rotation (that does not significantly change the Chandrasekhar mass), the disk mass is $M_D \\sim 10^{-2}-0.1 \\, M_\\odot$.  However, if the WD is uniformly rotating, the disk mass is significantly lower, even if the WD is rotating near break-up \\citep{Ott10}.  Most of the rotationally supported material is initially just outside the proto-neutron star at a radius of $R_D \\sim$ 100 km. The disk is born sufficiently hot that nuclei are dissociated into free nucleons and weak interaction processes have time to change the neutron to proton ratio. Under these conditions, the composition of the disk can be usefully described using the electron fraction \\begin{equation} Y_e = \\frac{n_p}{n_n+n_p} \\end{equation} where $n_p$ and $n_n$ are the proton and neutron densities, respectively. Weak interactions in the disk are dominated by the pair capture reactions \\begin{equation} e^+ + n \\rightarrow \\bar{\\nu}_{e} + p \\label{eq:p} \\end{equation} \\begin{equation} e^- + p \\rightarrow \\nu_{e} + n \\label{eq:n} \\end{equation} where the resulting neutrino emission is also the dominant coolant for the disk material because the photons are trapped and cannot escape on the timescales of interest.  Although the disk is formed very hot, electron degeneracy suppresses the formation of electron/positron pairs.  The resulting electron excess favors electron captures (eq.~[\\ref{eq:n}]) over positron captures (eq.~[\\ref{eq:p}]), thus driving the composition neutron rich (e.g.~\\citealt{Pruet+03}).  As the disk accretes onto the proto-neutron star, the remaining mass spreads to larger radii in order to conserve angular momentum.   As the disk spreads, the temperature drops (for the majority of the mass) and thus the neutrino cooling and weak reactions become less efficient. When the bulk of the disk reaches $R_D \\sim$ 1000 km, the disk has cooled to T $\\sim 10^{10}$K, and weak interactions freeze out \\citep{MPQ09-Neutron}.  Absent a central neutron star -- as is the case for analogous disks formed during binary neutron star  mergers -- weak interactions freeze out when the material is still quite neutron rich, with $Y_e \\sim 0.2-0.4$.   During AIC, however, the proto-neutron star emits copious neutrinos in the first few seconds after formation \\citep{BL}, which irradiate the disk as it spreads.  Because the neutron star is deleptonizing, electron neutrinos dominate the emission and the rate of neutrino captures $\\nu_{e} + n \\rightarrow p + e^{-}$ becomes comparable to the rate of pair captures (eqs.~[\\ref{eq:p}]-[\\ref{eq:n}]).  For a significant fraction of the mass, this increases the electron fraction at freeze-out to $Y_e \\gtrsim$ 0.5.  As the disk continues to spread, the free nucleons in the disk recombine to form heavy nuclei once T $\\sim$ few $\\times 10^9$ K.   The formation of He alone releases $\\sim 7$ MeV per nucleon, which is larger than the gravitational binding energy of the disk when fusion commences ($\\sim 2$ MeV per nucleon).  As a result, the disk is unbound, with its slow viscous spreading transformed into a supersonic outflow that reaches $v \\simeq 0.1$ c.  The end state of the fusion that initiates the disk's destruction is determined by the density, temperature, and electron fraction (the latter because weak interactions have frozen out); the parameters are such that the ejecta is nearly in nuclear statistical equilibrium (NSE). Since there are roughly equal numbers of protons and neutrons, a significant fraction of the unbound disk fuses to form $^{56}$Ni,  the decay of which can power a supernova-like transient once the ejecta expands sufficiently that photons may diffuse out.  In the remainder of this paper, we calculate the lightcurves and spectra of this transient.  \\vspace{-0.5cm}  \\label{sec:discussion}  Although the theoretical possibility of AIC has been recognized for decades (\\citealt{Canal&Schatzman76};\\citealt{Nomoto+79}), only recently have transient surveys had sufficient depth, field of view, and cadence to potentially detect and characterize these events.  The identification of AIC, or the ability to place meaningful constraints on its rate, would have important implications for a variety of astrophysical problems, including degenerate binary evolution, the progenitors of Type Ia SNe, and the formation channels of globular cluster NSs \\citep{Grindlay87}, low mass X-ray binaries \\citep{vandenHeuvel84}, and millisecond pulsars \\citep{Grindlay&Bailyn88}.  Because AIC transients are powered by the decay of $^{56}$Ni, they have many qualitative similarities to Type Ia SNe.  Although the observed lightcurves are too dim to perform cosmology, characterizing the correlation between peak luminosity and decay timescale (i.e., the \\citealt{Phillips93} relation) for AIC would be useful for observationally identifying these events.  For example, the duration and peak luminosity of AIC transients increases with the ejecta mass, which probably increases in proportion to the initial mass of the accretion disk \\citep{MPQ09-Nickel}.  The disk mass, in turn, is set by the angular momentum and rotational profile of the WD prior to collapse (\\citealt{Dess06}; \\citealt{Ott10}), which also determines whether AIC is expected to be a powerful source of gravitational waves (\\citealt{Ott09}).  Everything else being equal, more luminous AIC transients are thus more likely to be accompanied by a detectable gravitational waves, making AIC transients a promising electromagnetic counterpart to events detected with ground-based interferometers such as Advanced LIGO.  We now briefly address the prospects for the ``blind'' detection of AIC with current and future transient surveys.  Although the rate of AIC is uncertain, the amount of rare elements produced by the neutron-rich SN ejecta constrains the rate $R_{\\rm AIC}$ to be $\\lesssim 10^{-4}$ yr$^{-1}$ per Milky Way-type galaxy (or $\\lesssim 10^{-2}$ of the SN Ia rate) if the AIC calculations of \\citet{Dess06} are representative.  Assuming that the AIC rate is proportional to the blue stellar luminosity \\citep{Kopparapu+08}, a Galactic rate of $R_{\\rm AIC} \\equiv 10^{-4}R_{-4}$ yr$^{-1}$ corresponds to a volumetric rate of $10^{-6}R_{-4}$ Mpc$^{-3}$ yr$^{-1}$.  For an ejecta mass of $M_{\\rm ej} = 10^{-2}M_{\\sun}$, we find that AIC results in an optical transient with a peak luminosity $L_{\\rm peak} =$ few $\\times 10^{41}$ ergs s$^{-1}$ (Fig.~\\ref{fig:bol-lightcurves}).  The resulting luminosity distance is 600(400)[100]$L_{41}^{1/2}$ Mpc for a magnitude limit of 25(24)[21], where $L_{41} \\equiv L_{\\rm peak}/10^{41}$ ergs s$^{-1}$.  The PanSTARRs Medium Deep Survey (MDS; \\citealt{Kaiser+02}), which covers $\\sim 84$ deg$^{2}$ in $g$ and $r$ down to AB magnitude $\\sim 25$, should thus detect $\\sim 2 L_{41}^{3/2}R_{-4}$ AIC transients per yr.  The 5-day cadence survey of the Palomar Transient Factory (PTF; \\citealt{Law+09}), which surveys an active area $\\sim 2700$ deg$^{2}$ to a limiting AB magnitude of 21, will detect $0.3 L_{41}^{3/2}R_{-4}$ yr$^{-1}$.  Thus, for $R_{\\rm AIC} \\sim 10^{-4}$ yr$^{-1}$ and $L_{41} \\sim$ few we conclude that MDS and PTF should detect $\\sim 1-3$ AIC per year.  As we discuss below, identifying AIC transients for follow-up observations will be challenging, so these estimates are probably optimistic and motivate faster cadence surveys such as the Palomar P60-FasTING survey \\citep{Kasliwal+10}.  Prospects for detection are better with the Large Synoptic Survey Telescope (LSST; \\citealt{Strauss+:2010}), which will image the entire sky down to a limiting magnitude $\\sim 24.5$ every 3$-$4 nights and should detect AIC events at a rate $\\sim 10^3 \\, L_{41}^{3/2}R_{-4}$ yr$^{-1}$.    Unfortunately, the fraction of AIC that have ejecta masses of $\\sim 10^{-2} M_\\odot$, and thus peak luminosities of $L_{\\rm peak} =$ few $\\times 10^{41}$ ergs s$^{-1}$, is very difficult to predict given uncertainties in the differential rotation of the progenitor WD.   Absent additional theoretical progress, it will thus not be easy to turn upper limits from transient surveys into constraints on the rate of AIC.  Although the prospects that upcoming transient surveys will detect AIC are promising, their identification will be challenging due to the deluge of other transients that will be discovered (and, indeed, are already being discovered).  Other thermal transients predicted to occur on $\\sim$ few day timescales include the failed deflagrations of C/O WDs \\citep{Livne+05}, ``.Ia'' SNe due to unstable thermonuclear He flashes from WD binaries (\\citealt{Bildsten+07}; \\citealt{Shen+10}), shock breakout from Type 1a supernova \\citep{piro2010}, and radioactively-powered transients from NS mergers (\\citealt{Metzger10}).  Although NS mergers and other WD-related explosions may originate from a similar stellar population as AIC, a defining characteristic of AIC is that the majority of the ejected material is processed through NSE.  This means that few low or intermediate-mass elements are likely to be produced (such as He, O, C, and Ti), in contrast to what is expected from .Ia's or the failed detonation of a WD.  We note, however, that trace amounts of Ca and Ti are produced in some of our models (see Table \\ref{table:comp_disk}), which still may produce strong lines despite their low abundances (e.g.~the Ca triplet at $\\sim 9000$ \\AA; Fig.~\\ref{fig:color-lightcurves}).  A perhaps more dramatic difference between AIC transients produced by a disk around the newly formed NS and other WD-related events are the large outflow speeds $\\sim 0.1$c in AIC, which leads to highly velocity-broadened lines.  The speed of the ejecta from NS mergers, as in AIC, is expected to be large.  A distinguishing feature in this case, however, is that NS merger ejecta is more neutron-rich than AIC ejecta, because in AIC the accretion disk material that ultimately becomes unbound is first irradiated by electron neutrinos from the proto-neutron star (\\citealt{MPQ09-Neutron,MPQ09-Nickel}).  Although the current line data for the $r-$process nuclei produced in NS merger ejecta is insufficient to make concrete spectroscopic predictions, it is safe to say that NS mergers are unlikely to resemble any spectral class of SNe yet detected.  AIC and NS merger transients may also be distinguished by their light curve evolution.  In Figure \\ref{fig:aic-ns-bol-lightcurves} we compare the bolometric light curves from AIC (this paper) and NS mergers (from \\citealt{Metzger10}) assuming an ejecta mass $M_{\\rm ej} = 10^{-2}M_{\\sun}$ in both cases.  Although the overall brightness and late-time evolution is similar, the NS merger light curve shows an earlier, more pronounced peak.  This difference results from the different radioactive heating rate in the two cases.  AIC transients are powered by the decay of $^{56}$Ni, which produces an approximately constant heating rate prior to the $^{56}$Ni half-life $\\sim 6$ days.  NS merger transients, by contrast, are powered by a large number of decaying nuclei with a wide distribution of half-lives, resulting in a total heating rate that decreases approximately as a power-law $\\dot{E} \\propto t^{-1.2-1.4}$.   NS merger transients peak earlier simply because the radioactive heating is concentrated at early times.  Distinguishing between NS mergers and AIC based on light curve evolution alone would thus probably require catching the event prior to maximum.  \\begin{figure} \\centering \\begin{overpic}[scale=0.30]{lightcurves_with_NS_merger.ps} \\put(113,15){Time (d)} \\put(8,40){\\begin{sideways}Bolometric Luminosity $\\left(\\textrm{ergs s}^{-1}\\right)$\\end{sideways}} \\put(153,149){$Y_e : 0.425 - 0.55$} \\put(153,139.5){$Y_e : 0.45 - 0.55$} \\put(153,130.5){$\\epsilon_{\\textrm{therm}} = 0.5$} \\put(48,80){AIC} \\put(48,146){NS-NS merger} \\end{overpic} \\caption{Comparison of the bolometric light curves of the transients from AIC (this paper) and NS mergers \\citep{Metzger10}.  Both models were calculated for a total ejecta mass of $M_{\\rm tot} = 10^{-2}M_{\\sun}$.  The NS-NS merger transient is produced by the decay of a large number of neutron-rich nuclei created by r-process nucleosynthesis in the ejecta; \\citet{Metzger10} estimated a thermalization efficiency of $\\epsilon_{\\textrm{therm}} \\sim 0.5$ for these decays.  This figure highlights that distinguishing NS merger and AIC transients using lightcurve evolution alone will likely require early-time observations at $\\lesssim$ 1 day. \\label{fig:aic-ns-bol-lightcurves}} \\end{figure}  Several sub-luminous, rapidly-evolving Type I supernovae have recently been discovered, which could in principle be related to AIC.  These include SN 2008ha (\\citealt{Valenti+09}; \\citealt{Foley+09}; \\citealt{Foley+10}), 2005E (\\citealt{Perets+09}), and 2002bj (\\citealt{Poznanski+10}).  In all of these events, the low $^{56}$Ni masses inferred from their peak luminosities ($\\sim 10^{-3}-0.1M_{\\sun}$) are largely consistent with the quantity predicted from AIC.  However, the presence of low and intermediate mass elements (e.g. Ca, C, O, He) in the spectra of these events with relatively low velocities ($\\sim 2000-10000$ km s$^{-1}$) rules out the standard Ni-rich AIC scenario outlined in this paper.  As discussed by \\citet{MPQ09-Nickel}, the timescale, velocity, and composition of the AIC transient can be modified if a substantial quantity of WD material remains at large radii following AIC, either because the WD is still collapsing or because some matter remains centrifugally supported in a remnant disk created by a super-M$_{\\rm Ch}$ WD-WD merger (e.g. \\citealt{Yoon+07}).  In the latter case, depending on the mass and composition of the WDs, up to $\\sim 0.1M_{\\sun}$ in C, O, Ne, or He could remain at a radius $\\sim 10^{9}$ cm.  Once the SN explosion or Ni disk wind impacts the remnant material, it will (1) shock heat the material to a few times $10^{9}$ K, which may synthesize intermediate mass elements but will not disassociate the Ni and will leave some unburned WD material; (2) slow the ejecta from $v \\sim 0.1$ c to a few thousand km s$^{-1}$.  A combination of slower ejecta and higher opacity would lengthen the duration of the light curve, making such an event easier to detect.  Such an ``enshrouded AIC'' scenario may therefore still be capable of explaining events such as SN 2008ha, 2005E, and 2002bj.  In order to make a quantitative comparison to observations, however, hydrodynamic simulations will be required to determine the dynamics and nucleosynthetic yield of the complex interaction between the prompt SN, the Ni-rich disk wind, and the remnant WD material.   Moreover, the timescale of AIC relative to the WD-WD merger itself is uncertain; any significant delay would likely lead to less debris at large radii and a ``cleaner\" event analogous to that calculated in this paper.  On this note it is also important to distinguish between AIC due to accretion from a binary companion, as considered in this paper, and electron-capture supernova in AGB stars with O-Ne-Mg cores (for recent work, see \\citealt{Kitaura+06}; \\citealt{Wanajo+09}).  Although the collapse mechanism in these two cases is similar, an important difference in the AGB case is the presence of an extended hydrogen envelope.  This additional overlying material both increases the $^{56}$Ni yield of the explosion and extends the duration of the supernova due to its larger opacity (\\citealt{Wanajo+09}).  Although the lack of hydrogen in the spectra of SNe 2008ha, 2005E, and 2002bj rules out electron capture SNe in these events, electron capture events may explain slower evolving sub-luminous SN such as 2008S (e.g.~\\citealt{Prieto+08}; \\citealt{Thompson+09}; \\citealt{Pumo+09})  Our calculations could be improved in a number of ways.  In particular, future work will require more realistic multi-dimensional models in which the composition is calculated from a hydrodynamic simulation of the disk evolution beginning with free nucleons.   This is important for determining the final composition as a function of radius, which in turn can influence the spectroscopic predictions."
    },
    "1005/1005.1612_arXiv.txt": {
        "abstract": "We report an eclipse mapping study of the intermediate polar DQ~Her  based on time-resolved optical spectroscopy ($\\Delta\\lambda \\sim  3800-5000$\\,\\AA) covering 4 eclipses. The spectra were sliced into 295 narrow passbands in the continuum and in the lines, and the corresponding light curves were analysed to solve for a set of monochromatic maps of the disk brightness distribution and for the flux of an additional uneclipsed component in each band. Eclipse maps of the He\\,II $\\lambda 4686$ line indicate that an azimuthally- and vertically-extended bright spot at disk rim is important source of reprocessing of x-rays from the magnetic poles.   The disk spectrum is flat with no Balmer or Helium lines in the inner regions, and shows double-peaked emission lines in the intermediate and outer disk regions while the slope of the continuum becomes progressively redder with increasing radius. The inferred disk temperatures are in the range $T\\simeq 13500 - 5000\\,K$ and can be reasonably well described by a steady-state disk with mass accretion rate of $\\dot{M}=(2.7\\pm1.0)\\times10^{-9}\\,M_{\\odot}\\,yr^{-1}$. A comparison of the radial intensity distribution for the Balmer lines reveals a linear correlation between the slope of the distribution and the transition energy. The spectrum of the uneclipsed light is dominated by Balmer and He\\,I lines in emission (probably from the extended nova shell) with narrow absorption cores (likely from a collimated and optically thick wind from the accretion disk). The observed narrow and redshifted Ca\\,II $\\lambda 3934$ absorption line in the total light spectra plus the inverse P-Cygni profiles of the Balmer and He\\,II $\\lambda 4686$ emission lines in spectra of the asymmetric component indicate radial inflow of gas in the innermost disk regions and are best explained in terms of magnetically-controlled accretion inside the white dwarf magnetosphere. We infer projected radial inflow velocities of $\\sim200-500\\,km\\,s^{-1}$, significantly lower than both the rotational and the free-fall velocities for the corresponding range of radii.  A combined net emission He\\,II plus H$\\beta$ low-velocity eclipse map reveals a twisted dipole emitting pattern near disk center. This is interpreted as being the projection of accretion curtains onto the  orbital plane at two specific spin phases, as a consequence of the selection in velocity provided by the spectral eclipse mapping. ",
        "introduction": "In many cataclysmic variable stars, the magnetic field of the white dwarf (primary star) is sufficiently weak to be neglected. On the other hand, in systems like the AM Her stars, the magnetic field dominates completely the accretion flow \\citep{2001cvs..book.....H}. Intermediate cases can lead to a complex range of phenomena. This occurs in the intermediate polar stars, where a magnetic field of moderate strength ($B \\lesssim 10^{7}\\,G$) allows the combination of characteristics of non-magnetic systems (in the outer disk regions) with characteristics of highly magnetic systems (in the regions close to the white dwarf).  In intermediate polars the inner regions of the accretion disk are disrupted by the magnetic field of the white dwarf and the infalling plasma is forced to follow the field lines down to the star's surface at one or both the magnetic poles, transforming kinetic energy into radiation. Part of this energy can be reprocessed in other sites of the binary, leading to pulsations at the white dwarf rotation period and/or its beating with the orbital period \\citep{1986ApSS.118..271W}.  DQ Her, or Nova Her 1934, is a deeply eclipsing cataclysmic variable (orbital period $P=4.6\\,h$) and the prototype of the subclass of intermediate polars, showing a coherent $71\\,s$ pulsation \\citep{1956ApJ...123...68W}. Due to the high inclination of the system \\citep[$i=86.5\\degr$;][]{1993ApJ...410..357H} the white dwarf, accretion disk and bright spot are occulted by the secondary star, allowing the emission from these different light sources to be separated and spatially resolved by indirect imaging techniques such as eclipse-mapping \\citep{1985MNRAS.213..129H} and Doppler tomography \\citep{1988MNRAS.235..269M}.  Here we report the results of an eclipse-mapping study of the spectra and structure of the accretion disk of DQ~Her based on time-resolved spectroscopy of 4 eclipses collected on July 1987. The observations are described in Section~\\ref{obs}. Data analysis procedures are given in Section~\\ref{data_ana}. In Section~\\ref{results} we investigate the disk spatial structure in the emission lines and in the continuum and we present spatially resolved spectra of the disk, gas stream and of the uneclipsed light. The results are discussed in Section~\\ref{discussion} and summarized in Section~\\ref{summ}. In the Appendix~\\ref{appendixA} we discuss the ability to reconstruct the full-width-half-maximum distribution of emission lines with spectral mapping techniques. ",
        "conclusions": "\\label{discussion}  \\subsection{An energy dependency of the emission-line radial intensity distribution slope} \\label{energy-slope}  In Section~\\ref{emission} we found that the slope of the radial line intensity distributions correlates with the transition energy of the Balmer lines.  In order to extend our analysis we extracted (using the same procedures described in Section~\\ref{emission}) the radial intensity distribution for H$\\epsilon$, the next line in the Balmer series. We then fitted a power-law $I \\propto r^{-\\alpha}$ \\citep[e.g.,][]{1981AcA....31..395S, 1986MNRAS.218..761H} for each radial distribution in order to find the best-fit $\\alpha$ index for each line. The innermost disk regions ($R < 0.2~R_{L1}$) are excluded from the fit since the continuum reaches the emission level of the lines in the regions inside this radius. The results are listed in Table~1 and plotted in Fig.~\\ref{alphas}.  \\begin{table} \\begin{center} \\caption{Index $\\alpha$ for the Balmer series and He\\,II lines.} \\vspace{12 pt} \\begin{tabular}{l c c c c} \\hline \\hline Line        & E[eV] & $\\alpha$ ($I \\propto r^{-\\alpha}$) \\\\ \\hline H$\\beta$    & 12.76 &  $1.20 \\pm 0.06$ \\\\ H$\\gamma$   & 13.06 &  $2.16 \\pm 0.11$ \\\\ H$\\delta$   & 13.23 &  $2.62 \\pm 0.10$ \\\\ H$\\epsilon$ & 13.33 &  $2.95 \\pm 0.12$ \\\\ He\\,II 4686 & 75.50 &  $3.15 \\pm 0.14$ \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\begin{figure} \\includegraphics[bb=1.5cm -7cm 20cm 18cm,angle=-90,scale=0.40]{fig10.eps} \\caption{Distribution of the index $\\alpha$ in function of the transition energy for the He\\,II 4686 and Balmer series lines. The solid line marks a linear fit for the Balmer lines distribution (see Section~\\ref{energy-slope}).} \\label{alphas} \\end{figure}  The $\\alpha$ index increases with the transition energy for the Balmer lines, i.e., the emission at lines with larger transition energy (e.g., H$\\epsilon$) are more concentrated at the disk center than lines of lower energy (e.g., H$\\beta$). Moreover, our results reveal that there is a strong linear correlation between $\\alpha$ and the transition energy for the Balmer lines (see Fig.~\\ref{alphas}). This result is in contrast with the empirical $I \\propto r^{-1.5}$ law \\citep{1990ApJ...364..637M} usually evoked in works about radial distribution of emission lines in accretion disks. The He\\,II 4686 line radial distribution falls much more rapidly with radius than the Balmer lines, with an $\\alpha = 3.15$, and does not fit the linear relation of the Balmer lines. This is not surprising, as significantly different physical conditions are required to produce the H\\,I and He\\,II lines.  \\cite{1969AcA....19..155S} and \\cite{1972ApJ...171..549H} showed that the double-peaked profile of emission lines in accretion disks can be accounted for by making two basics assumption: ($i$) the geometry is flat, i.e., the disk emission is restricted to the orbital plane, and ($ii$) the disk is in Keplerian motion. \\cite{1981AcA....31..395S} showed that the line profile is sensitive to the line radial intensity distribution and that the $\\alpha$ value could be recovered by fitting a model to the line profile.  From observational evidences \\citep[e.g.,][]{1979ctvs.conf...65F,1980ApJ...238..955Y} he found that the emission lines profiles could be represented by power-laws $I \\propto r^{-\\alpha}$, where $\\alpha$ is expected to be in the range $1-2$.  \\cite{1981AcA....31..395S} gives a list of results where $\\alpha$ could be estimated for Helium and Balmer series lines (Z~Cha, \\citealt{1979ctvs.conf...65F}; SS~Cyg, \\citealt{1980ApJ...241..269C,1980ApJ...240..597S}; HT~Cas, \\citealt{1981ApJ...245.1035Y}; DQ~Her, \\citealt{1979ApJ...232..500H,1980AZh....57..749D}). In all cases the value obtained for $\\alpha$ is in the range $1.7-2.2$ suggesting that there are no systematic differences between objects of different classes of cataclysmic variables. Similar values of $\\alpha$ are obtained by fitting the line profile of the dwarf novae U~Gem or Z~Cha and the intermediate polar DQ~Her. \\cite{1979ApJ...232..500H} found that in DQ~Her the H$\\beta$ line has $\\alpha=2.2$ while He\\,II 4686 has $\\alpha=1.9$.  These values are inconsistent with our results. The discrepancy may be explained by the low spectral resolution of the data of \\cite{1979ApJ...232..500H}, which obtained $\\alpha$ from line profiles of low resolution (dispersion of 48 \\AA\\, mm$^{-1}$) while we extracted $\\alpha$ from spatially resolved spectra obtained from data of significantly higher spectral resolution (2 \\AA\\,pixel$^{-1}$).  \\cite{1987MNRAS.228..779M} analysed the Balmer lines in the dwarf nova Z Cha. In his study the differences between the line profiles in the average spectrum are evident. The line wings become wider from H$\\alpha$ to H$\\delta$ indicating that the radial line distributions become more concentrated for increasing transition energy. In quiescence, Z Cha has a relatively small inferred mass transfer rate of $\\dot{M}=10^{-12}\\,M_{\\odot}\\, yr^{-1}$. Simulations shows that the LTE model is not valid in this case, producing a poor fit to the Balmer lines \\citep{1987MNRAS.228..779M}.  \\cite{1988MNRAS.231.1117M} found $\\alpha=1.8$ for the Balmer lines in the dwarf nova IP~Peg in quiescence. Similarly, \\cite{1990ApJ...349..593M} showed that in IP~Peg the radial distribution of the Balmer lines have similar slopes when the object is in quiescence but that the radial distributions are different for distinct lines when the object is in outburst. The H$\\delta$ line distribution is steeper than H$\\gamma$, that is steeper than H$\\beta$, i.e., the $\\alpha$ index of the radial profile increases with the transition energy. A similar result was found by \\cite{1990ApJ...364..637M} for the dwarf nova U~Gem, with H$\\gamma$ and H$\\beta$ showing, respectively, $\\alpha=1.66$ e $1.43$. In contrast, for the white dwarf V2051~Oph, \\cite{2006AJ....131.2185S} found $\\alpha=1.55$ for H$\\delta$ and $\\alpha=1.78$ for H$\\gamma$; $\\alpha$ decreases with the transition energy in this dwarf nova. During their observations V2051~Oph was in an unusual low brightness state (low mass transfer rate). Thus, there are observational evidences that the slope of the radial line distribution is different for different emission lines.  In order to understand the physics behind the line formation in accretion disks it is necessary to combine new and improved theory with spatially-resolved observational studies. Spatially resolved studies from Doppler tomography and eclipse mapping techniques can provide high resolution information about the line distribution in the accretion disk. More sophisticated spectral models for accretion disk need to include features such as the vertical temperature structure and external sources of heat.   \\subsection{Are we seeing the accretion curtains?} \\label{curtains}  Here we discuss possible interpretations for the asymmetries seen in the low-velocity eclipse maps of the H$\\beta$ and He\\,II $\\lambda 4686$ lines. These asymmetries cannot be explained in terms of gas stream overflow \\citep[e.g.,][see their Fig.~5]{2004AJ....128..411B} as they are  physically unrelated to the gas stream trajectory, nor can they be interpreted in terms of spiral shock waves \\citep[e.g.,][see their Fig.~2]{2000MNRAS.314..727B} as the morphology, orientation and location are inconsistent with those expected for two-armed  tidally-induced spiral structures. Moreover, it is also not possible to explain the observed structures in terms of a front-back asymmetry caused by a flared disk \\citep[e.g.,][see their Figs.~3 and 8]{1998MNRAS.298.1079B} because this would affect both line and continuum maps, whereas in DQ~Her the asymmetries are only seen at specific velocity-resolved eclipse maps of particular emission lines. Additionally, the disk opening angle is not large enough to result in measurable obscuration or limb-darkening effects (see Section~\\ref{mem}).  The asymmetries observed in the H$\\beta$ and He\\,II $\\lambda 4686$ line maps seem to be of a different nature, related to something not previously seen in eclipse mapping experiments.  DQ~Her is the first intermediate polar being eclipse mapped. It is tempting to interpret the structures seen in the low-velocity He\\,II and H{$\\beta$} maps in terms of the projection onto the orbital plane of accretion columns/curtains formed over the magnetic poles of the white dwarf by the channeled material \\citep[e.g.,][]{1996PASA...13...87F}\\footnote {The eclipse mapping procedure has the approximate effect  of projecting any emission from above the orbital plane back along the  inclined line of sight to the point where it pierces the disk plane.  Thus, emission arising at height $Z$ above the disk is projected a  distance $\\rm Z \\, {\\it tan}\\, \\rm i$ towards the back of the disk and is distorted by the intrinsic azimuthal smearing effect of the eclipse mapping method \\citep{1995ApJ...448..395B}. With our choice of azimuthal blur width $\\Delta\\phi$ (Section~\\ref{mem}), any point source in the eclipse maps will appear smeared in azimuth by $\\simeq 20\\degr$.}. In this scenario, due to different opacities for incoming irradiating photons, and energies needed to ionize atoms and to excite electrons to the upper energy levels of each transition, the regions of the accretion curtains close to the magnetic poles are prone to reprocess the strong irradiation by hard x-rays photons whereas regions farther away receive less energetic photons. Hence, high-excitation lines such as He\\,II $\\lambda 4686$ trace the emission in the accretion curtains regions closer to the magnetic poles, whereas  the low-excitation Balmer lines trace emission from regions farther away from the x-ray emitting magnetic poles. Lines of different excitation frame different parts of the accretion flow and, individually, yield only a partial view of the emitting accretion curtains. According to this, a better view of the accretion curtains may be obtained by combining the complementary high-excitation He\\,II map with the low-excitation H$\\beta$ map. An eclipse map obtained by combining the net line emission He\\,II $\\lambda 4686$ and H$\\beta$ low-velocity eclipse maps is shown in  Fig.~\\ref{disk_soma}. A remarkable structure resembling a twisted dipole emitting pattern can be seen near disk center. Although this figure encourages one to further pursue this line of reasoning, it poses an immediate problem.  \\begin{figure} \\includegraphics[bb=1cm -15cm 20cm 18cm,angle=-90,scale=.3]{fig11.eps} \\caption{Combined net emission He\\,II plus H$\\beta$ low-velocity eclipse map. The notation is similar to that of Fig.~\\ref{hbet_curves}.} \\label{disk_soma} \\end{figure}  Each eclipse map yields a snapshot of the disk brightness distribution averaged over the period covered by the 4 eclipses. Therefore, because it comprises a timescale much longer than the 71\\,$s$ spin period of the white dwarf, changes in brightness distribution caused by the varying aspect of the accretion curtains with the white dwarf rotation should be completely smeared out in the eclipse map --- just as in a long exposure picture of a moving target. How could one then obtain a coherent image of fast rotating accretion curtains in an eclipse map? The answer can be built in terms of the selection in velocity for the accreting material performed by the application of spectral mapping techniques.  The velocity field for gas in the accretion curtains is a sum of its inflowing component (along the magnetic field lines) and a rotational  component (due to the white dwarf spin period). As seen in Sections~\\ref{spectra}~and~\\ref{emission}, the velocity field is dominated by the rotational component, which is significantly larger than the  inflowing component. At most rotation phases (i.e., curtain orientations), different parts of the curtains show distinct projected velocities along the line of sight and their emission will split onto several different velocity bins across the emission line profile.  Conversely, a specific velocity-resolved eclipse map will contain a mixture of different parts of the accretion curtains at different rotation phases. This slicing and mixing effect prevents one from obtaining a coherent image of the accretion curtains at all velocity-resolved eclipse maps, with the fortunate exception of the low-velocity bin.  When the curtains are aligned with the binary axis --- and only at this orientation --- the rotational component will be suppressed (because it becomes orthogonal to the line of sight), and the curtains will show only the small inflowing velocity component (see Section~\\ref{spectra}), contributing to the same particular low-velocity bin. We note that two effects contribute to further reduce the observed projected velocity: (i) the inflow velocity must decrease after the accretion shock until it reaches zero at the white dwarf surface, and (ii) for accretion curtains twisted off the orbital plane, regions close to the magnetic poles present the lowest projected velocity along the line of sight. Therefore, looking at a low velocity bin is equivalent of framing the accretion curtains always at the same two spin phases (where the curtains are oriented at an azimuth roughly aligned with the binary main axis), allowing one to obtain a coherent image of them.  A further question is why the accretion curtains lag the binary main axis by $\\simeq 30\\degr$. \\cite{1999ApJ...519..324H} applied Doppler tomography techniques to spin-pulse spectroscopy of intermediate polars to find that the optical pulsation arises from azimuthally-extended accretion curtains, located typically a few white dwarf radii from disk center. His tomogram of the slow rotator PQ~Gem shows a twist in the accretion curtains, with material feeding onto magnetic field lines $\\sim 40\\degr$ ahead of the location at which it hits the white dwarf. This is the opposite case of what is seen in DQ~Her and the difference may be in the spin period. For comparable magnetic field strengths and mass accretion rates a slow rotator (the case of PQ~Gem, $P_{spin}=834\\,s$) picks material at a magnetospheric radius with higher (Keplerian) velocity than the accretion curtains. Thus, the accretion curtains tend to step ahead of the magnetic pole. On the other hand, in a fast rotator (the case of DQ Her, $P_{spin}=71\\,s$) the field lines at the magnetosphere are rotating faster than the disk gas being channelled onto the field lines (see Section~\\ref{emission}). In this case the accretion curtains tends to lag the magnetic pole."
    },
    "1005/1005.4201_arXiv.txt": {
        "abstract": "The stability of an Einstein static universe in the DGP braneworld scenario is studied in this letter. Two separate branches denoted by $\\epsilon=\\pm1$ of  the DGP model are analyzed. Assuming the existence of a perfect fluid with a constant equation of state, $w$, in the universe, we find that, for the branch with $\\epsilon=1$, there is no a stable Einstein static solution, while, for the case with $\\epsilon=-1$, the Einstein static universe exists and it is stable when $-1<w<-\\frac{1}{3}$.  Thus, the universe can stay at this stable state past-eternally and may undergo a series of infinite, non-singular oscillations.   Therefore, the big bang singularity problem in the standard cosmological model can be resolved. ",
        "introduction": "Although most of the problems in the standard cosmological model can be resolved by the inflation theory, the resolution of the existence of a big bang singularity in the early universe is still elusive. Based upon the string/M-theory, the pre-big bang~\\cite{Gasperini2003}  and ekpyrotic/cyclic~\\cite{Khoury2001} scenarios have been proposed to address the issue. Recently, Ellis et al proposed, in the context of general relativity, a new scenario, called an emergent universe~\\cite{Ellis2004a, Ellis2004b} to avoid this singularity. In this scenario, the space curvature is positive, which is supported by the recent observation from WMAP7~\\cite{Komatsu2010} where it was found that a closed universe is favored  at the $68\\%$ confidence level,  and the universe stays, past-eternally, in an Einstein static state and then evolves to a subsequent inflationary phase. So, in an emergent theory, the universe originates from an Einstein static state rather than from a big bang singularity.  However, the Einstein static universe in the classical general relativity is unstable, which means that it is extremely difficult for  the universe  to remain in such an initial static state in a long time due to the existence of perturbations, such as the quantum fluctuations. Therefore,  the original emergent model does not seem to resolve the big bang singularity problem successfully as expected.  However, in the early epoch, the universe is presumably under extreme physical conditions, the  realization of the initial state may be affected by  novel physical effects, such as those stemming from quantization of gravity, or a modification of general relativity or even other new physics. As a result, the stability of the Einstein static state has been examined in various cases~\\cite{Carneiro2009, Mulryne2005, Parisi2007, Wu2009, Lidsey2006, Bohmer2007, Seahra2009, Bohmer2009, Barrow2003, Barrow2009, Clifton2005,Boehmer2010, Boehmer20093, Wu20092, Odrzywolek2009},  from loop quantum gravity ~\\cite{Mulryne2005, Parisi2007, Wu2009} to  modified gravity (for a review see Ref.~\\cite{Boehmer2010} ), from Horava-Lifshitz gravity~\\cite{Boehmer20093, Wu20092} to Shtanov-Sahni braneworld scenario~\\cite{Lidsey2006}      In this paper,  we plan to examine the stability of the Einstein static universe in the DGP brane-world model~\\cite{Dvali2000}. In this braneworld, the whole energy-momentum is confined on a three dimensional brane embedded in a five-dimensional, infinite-volume Minkowski bulk. Since there are two different ways to embed the 4-dimensional brane  into the 5-dimensional spacetime, the DGP model has two separate branches denoted by $\\epsilon$ with distinct features. The $\\epsilon=+1$ branch  can explain the present accelerating cosmic expansion without the introduction of dark energy~\\cite{Deffayer2001}, while for the $\\epsilon=-1$ branch, dark energy is needed in order to yield an accelerating expansion, as is the case in the LDGP model~\\cite{Lue2004} and the QDGP model~\\cite{Chimento2006}. Using the $H(z)$, CMB shift and Sne Ia observational data,  Lazkoz and Majerotto~\\cite{Lazkoz2007} found that the LDGP and QDGP are slightly more favored than the self-accelerating DGP model. Let us also note that a crossing of a phantom divide line, which is favored by the recent various observational data~\\cite{Alam2004}, is possible with a single scalar field~\\cite{Zhang20062, Chimento2006}  in the $\\epsilon=-1$ branch. In addition, inflation in the DGP  model displays some new characteristics. It should be noted, however, that only in the $\\epsilon=-1$ case can inflation exit spontaneously~\\cite{Bouhmadi-Lopez2004, Cai2004, Papantonopoulos2004, Zhang2004, Zhang2006, Campo2007}. Also, in contrast to the Randall-Sundrum~\\cite{Randall1999} and Shtanov-Sahni~\\cite{Shtanov2003} braneworld scenarios with high energy modifications to general relativity, the DGP brane produces a low energy modification (for a review of the phenomenology of the DGP model, see Ref.~\\cite{Lue2006}). ",
        "conclusions": "In this paper, we have studied the existence and stability of the Einstein static universe in the DGP braneworld scenario. By assuming the existence of a perfect fluid with a constant equation of state, which is a good approximation if the perfect fluid is a scalar field and the variation of the potential of scalar field is very slow with time, we have shown that for the branch with $\\epsilon=1$, there is no stable Einstein static universe, whereas, for the branch with $\\epsilon=-1$, the Einstein static universe exists and it is stable if the equation of state $w$ satisfies $-1<\\omega<-\\frac{1}{3}$. Thus, the universe can stay at this stable state past-eternally, and may undergo a series of infinite, non-singular oscillations. Hence, in the $\\epsilon=-1$ branch of the DGP model, the universe can originate from an Einstein static state and then enter an inflation era. Furthermore, the universe can exit, spontaneously,  this inflation phase to a radiation dominated era, as shown in previous studies~\\cite{Bouhmadi-Lopez2004, Cai2004, Papantonopoulos2004, Zhang2004, Zhang2006, Campo2007}. As a result, the big bang singularity problem in the standard cosmological scenario can be resolved successfully."
    },
    "1005/1005.0564_arXiv.txt": {
        "abstract": "{} {We present the results of contemporaneous spectroscopic and photometric monitoring of the young solar-type star \\object{HD~171488} ($P_{\\rm rot}\\simeq 1.337$ days) aimed at studying surface inhomogeneities at both photospheric and chromospheric levels.} {Echelle FOCES spectra ($R\\simeq$\\,40\\,000) and Johnson $BV$ photometry have been performed in August 2006, with a good coverage of rotational phases. Spectral type, rotational velocity, metallicity, and gravity were determined using a code developed by us (\\textsc{Rotfit}) and a library of spectra at the same resolution of slowly-rotating reference stars.  The metallicity was measured from the analysis of iron lines with the \\textsc{Moog} code. The spectral subtraction technique was applied to the most relevant chromospheric diagnostics included in the FOCES spectral range, namely \\ion{Ca}{ii} IRT, H$\\alpha$, \\ion{He}{i}\\,D$_3$, H$\\beta$, and \\ion{Ca}{ii} H\\&K lines. } {A simple model with two large high-latitude spots is sufficient to reproduce the $B$ and $V$ light curves as well the radial velocity modulation, if a temperature difference between photosphere and spots of about 1500\\,K is used. A Doppler imaging analysis of photospheric lines basically confirms a similar spot distribution. With the help of an analogous geometric two-spot model, we are able to reproduce the observed modulations in the residual chromospheric emissions adopting different values of ratios between the flux of plages and quiet chromosphere (about 5 for H$\\alpha$ and 3 for \\ion{Ca}{ii} diagnostics). Facular regions of solar type appear to be the main responsible for the modulations of chromospheric diagnostics. Both the spot/plage model and the cross-correlation between the light curve and the chromospheric line fluxes display a sensible lead effect of plages with respect to spots (from 20$\\degr$ to 40$\\degr$ in longitude), as already observed in some active solar-type stars and RS CVn systems.} {The contemporaneous monitoring of photospheric and chromospheric diagnostics in the young and rapidly rotating solar-type star HD~171488 allowed us to detect active regions which have nearly the same location at both atmospheric layers, with plages slightly leading spots in longitudes. These active regions are similar to the solar ones in some respect, because the spot temperature is close to that of sunspot umbrae and the plage flux-contrast is consistent with the average solar values. The main differences with respect to the Sun are larger sizes and higher latitudes. } ",
        "introduction": "\\label{sec:Intro}  Young solar-mass stars just arrived on the Zero-Age Main Sequence (ZAMS) or on their way to it are in a very important phase of their life. Indeed, at this time they start to spin up approaching the ZAMS when they get free from their circumstellar disks. Thereafter, the disks can start to ``condense\" giving rise to proto-planetary systems and the stars experience angular momentum loss resulting from a magnetized stellar wind. It is therefore very important to define the physical conditions that affect the subsequent evolution of the star and of its environment. In addition to the basic stellar parameters ($T_{\\rm eff}$, $\\log g$, [Fe/H]), it is of fundamental importance to know the rotation rate and the level and behaviour of the star's magnetic activity.  Fast rotators with age of about 100 Myr are likely originated by a less efficient disk locking during their T~Tau phase, but they also suggest  a reduced effect of magnetic braking. It has been proposed that the rate of angular momentum loss decreases or, at least, saturates at high rotation rates (\\citealt{MacGregor91}; \\citealt{Barnes96}). Saturation of magnetic activity is witnessed by the typical behaviour of X-ray emission and other activity signatures as a function of the rotation speed (\\citealt{Randich97}; \\citealt{Krishna97}). Periodic variability of saturated fast-rotating stars is present both in the optical continuum (e.g., \\citealt{Marilli97}; \\citealt{Stassun04}) and in the X-ray band (\\citealt{Flaccomio05}). This rises questions about the true extent of saturation in fast-rotating stars. It is intriguing that very fast rotating stars, for which a ``super saturation\" regime, i.e. a slight decline of magnetic activity with the increase of rotation rate, is displayed  \\citep{Prosser96}, still show organized magnetic topologies but inefficient angular momentum losses \\citep{Stepien01}.  For all these reasons we consider it mandatory to investigate the surface topology of magnetic active regions on fast-rotating young stars, also because of the implications on the stability and early evolution of planetary systems.  HD~171488 (=\\object{V889~Her}) is a young fast-rotating solar analogue that has been the subject of several studies in the last decade. Its kinematic properties indicate that HD~171488 is a member of the Local Association, a stream of young stars with ages ranging from 20 to 150 Myr \\citep{Montes01}. It is probably the brightest single early-G type star that is rotating fast enough for mapping its photosphere through Doppler Imaging (DI). \\citet{Strass03} made the first dedicated photometric and spectroscopic study of HD~171488 in which they defined its astrophysical parameters, concluding that it is a single G0V main-sequence star with an age of about 30--50 Myr based on  both its position on the HR diagram and the lithium content. They also made the first reconstruction of its photosphere through the DI technique, showing a big polar spot with additional high-latitude features. Subsequent works based on DI have always found a polar spot \\citep{Marsden06,Jeffers08}. From long-term photometry and DI, two active longitudes separated by about 180$^{\\circ}$ seem to be persistent for several years \\citep{Jarvi08,Huber09}. \\citet{Jeffers08} found a high solar-type differential rotation with the equator lapping the poles every 12-13 days. \\citet{Jarvi08} found instead a much weaker differential rotation.  Despite the number of works devoted to this object in the last few years, which are mainly based on DI or Zeeman-DI, a simultaneous detailed investigation of the chromospheric and photospheric inhomogeneities has not been performed so far.  In the present work we analyze contemporaneous $BV$ photometry and high-resolution spectra of HD~171488. We determined the main parameters of the photospheric spots by means of the DI technique applied to a few suitable absorption lines and of a spot model for simultaneous solution of light and radial velocity curves. The behaviour of chromospheric inhomogeneities was investigated thanks to the Ca\\,{\\sc ii} H\\,\\&\\,K ($\\lambda=$3968.49 \\AA, 3933.68 \\AA), H$\\epsilon$ ($\\lambda=$ 3970.074 \\AA), H$\\alpha$ ($\\lambda=$6562.849 \\AA), and Ca\\,{\\sc ii} infrared triplet (IRT; $\\lambda=$8498.06 \\AA, 8542.14 \\AA, 8662.17 \\AA) lines.  The work is organized as follows. In Section\\,\\ref{sec:Obs} we describe the observations and data reduction. The spectral classification, rotational and radial velocities, evolutionary status, and metallicity are discussed in Section\\,\\ref{sec:AP}. The diagnostics of photospheric and chromospheric activity are analyzed in Sections \\ref{sec:phot_act} and \\ref{sec:chrom_act}, respectively.  A short discussion on the behaviour of photospheric and chromospheric activity is presented in Section\\,\\ref{sec:discussion}. Section \\ref{sec:concl} contains our conclusions. ",
        "conclusions": "\\label{sec:concl}  We performed an accurate analysis of high-resolution spectra and contemporaneous $BV$ photometry of the nearby young solar-type star HD~171488 with the main aim to characterize the behaviour of spots and plages on its surface. Revised astrophysical parameters, such as $T_{\\rm eff}$, $\\log g$, [Fe/H], rotational and heliocentric radial velocity, and lithium abundance, were also determined and have been found in close agreement with previous determinations from the literature, with the exception of the metallicity for which we did not find a sub-solar abundance as suggested by \\citet{Strass03}. Both the position on the HR diagram and the lithium abundance give an estimate of about 50\\,Myr for the age of HD~171488, confirming it as a young Sun on its way to the ZAMS.  We were able to find a clear radial velocity modulation  with an amplitude of about 500\\,m\\,s$^{-1}$ due to a ``Rossiter-McLaughlin'' effect caused by starspots. A spot modeling of the $B$ and $V$ light curves, performed with our code \\textsc{Macula}, is also able to reproduce the radial velocity modulation by adopting only two high-latitude spots 1500\\,K cooler than the quiet photosphere. A more accurate surface reconstruction was also performed by applying the Doppler Imaging analysis to two photospheric spectral lines and the $B$ and $V$ light curves. The method provides maps that essentially confirm the large polar spotted area with a temperature difference $\\Delta T\\simeq$\\,1500\\,K with respect to the unspotted photosphere.  Chromospheric activity was also studied with different spectral diagnostics. Equivalent widths of the residual emissions in the cores of H$\\alpha$, \\ion{Ca}{ii} H\\&K, and \\ion{Ca}{ii} IRT lines show clear modulations with the rotational phase and are ascribed to the presence of active chromospheric regions analogue to solar plages. A small but significant emission in the H$\\beta$ core, as well as an absorption feature corresponding to the \\ion{He}{i} $\\lambda5876$ line, are emphasized by the subtraction of the non-active template and are additional proofs of the high activity level of HD~171488 also in the upper chromosphere. Due to the fairly large relative errors, no clear rotational modulation could be detected in the H$\\beta$ and \\ion{He}{i} line fluxes. The average Balmer decrement, measured as the ratio of H$\\alpha$ to H$\\beta$ flux, confirms that the observed chromospheric emission is mainly coming from extended facular regions.  The H$\\alpha$ and \\ion{Ca}{ii} flux curves are well fitted by a model with two circular bright spots (plages), also performed with \\textsc{Macula}, with sizes similar to the underlying spots and a flux contrast of 5 for the H$\\alpha$ and about 3 for the \\ion{Ca}{ii} H\\&K and IRT fluxes, as expected from the different modulation amplitude. The flux curves are nearly anticorrelated with the photometry, even if the plage at a lower latitude precedes in longitude, by about 20$\\degr$--40$\\degr$, the corresponding photospheric spot. This finding is not anomalous since lead and lag between plages and spots have been observed in the Sun and are frequently invoked as the cause of phase shifts between light curves and chromospheric line fluxes observed both in young active stars and in RS~CVn systems."
    },
    "1005/1005.5388_arXiv.txt": {
        "abstract": "We use analytical and N-body methods to examine the survival of wide stellar binaries against repeated encounters with dark substructures orbiting in the dark matter haloes of dwarf spheroidal galaxies (dSphs). Our models adopt cosmologically-motivated conditions wherein dSphs are dark-matter dominated systems that form hierarchically and orbit about a host galaxy. Our analytical estimates show that wide binaries are disrupted at a rate that is proportional to the local density of dark substructures averaged over the life-time of the binary population. The fact that external tides can efficiently strip dark substructures from the outskirts of dSphs implies that the present number and distribution of binaries is strongly coupled with the mass evolution of individual galaxies. Yet we show that for the range of dynamical masses and Galactocentric distances spanned by Milky Way dSphs, a truncation in the separation function at $a_{\\rm max}\\simless 0.1 $pc is expected in {\\it all} these galaxies. An exception may be the Sagittarius dSph, which has lost most of is dark matter envelope to tides and is close to full disruption. Our simulations indicate that at separations larger than $a_{\\rm max}$ the perturbed binary distribution scales as $dN/da\\propto a^{-2.1}$ independently of the mass and density of substructures. These results may be used to determine whether the binary separation function found in dwarf galaxies is compatible with the scale-free hierarchical picture that envisions the existence of dark substructures in all galactic haloes. We show that the ACS camera on board of the Hubble telescope may be able to test this prediction in dSphs at heliocentric distances $\\simless 100$ kpc, even if the binary fraction amounts only 10\\% of the stellar population. ",
        "introduction": "\\label{sec:intro} Perhaps one of the most striking predictions from the present cosmological paradigm (Cold Dark Matter, CDM) is the existence of dark matter haloes devoid of baryonic material, the so-called ``dark galaxies''. Such systems arise because primordial gas is unable to cool efficiently by collisional exicitation of $H_2$ molecules (Haiman et al. 2000), whose formation is limited to haloes with virial temperatures above $10^4 $K (i.e. virial masses $\\simgreat 10^8 M_\\odot$). Haloes with masses below $\\sim 10^8 M_\\odot$ are thus expected to be dark. The fact that the CDM halo mass function diverges as $dN/d m\\propto m^{-1.9}$ (Diemand et al. 2008; Springel et al. 2008 and references therein) implies that dark galaxies dominate in number over luminous ones. Indeed, we expect of the order of $10^{12}$ subhaloes with masses within $10^{-6}$--$10^8 M_\\odot$ lingering in a Milky Way-size halo (Diemand et al. 2005).   How can we possibly test the existence of dark galaxies? If dark matter is in the form of a neutral relic particle once in thermal equilibrium in the early universe, these objects may be lit up by pair annihilation that would result in gamma-ray emmissions. Unfortunately, testing this scenario on the mass scale of dark galaxies is well beyond the present technological capabilities for the vast majority of particle candidates (e.g. Diemand et al. 2007). The second alternative is to detect dark haloes via gravitational perturbations of visible objects.  In this regard very wide binaries, i.e. those with separations $a\\simgreat 100$ AU, may enable simple experiments to test for the presence of dark substructures, for even extremely weak tidal perturbations can disrupt them. In the stellar halo of the Milky Way (hereafter MW), perturbers can be identified as inhomogeneities in the Galactic potential. Some examples are stellar clusters, molecular clouds and compact objects (i.e. MACHOS), which altogether cover a mass range that overlaps with that of dark galaxies (Carr \\& Sakellariadou 1999, Chanam{\\'e} \\& Gould 2004). Hence isolating the effects of different pertubers on binaries is not straightforward.  Much more interesting targets are binaries in {\\it dwarf galaxies}. The reasons are several: (i) dwarf spheroidal galaxies (hereafter dSphs) have the largest dark matter content of any known galaxy type in the Universe (e.g. Mateo 1998, Gilmore et al. 2007). In dSphs the fraction of dark to baryonic mass is so large ($10^1$--$10^3$) that the presence of baryonic inhomogeneities in the dwarf potential can be safely neglected. (ii) As a result of their hierarchical formation, dwarf galaxies orbiting the MW are expected to contain bound substructures. These objects are typically referred to as (sub-)subhaloes (hereafter sub$^2$haloes) in cosmological simulations. Given that the estimated virial mass of the MW dSphs is of the order of $M_{\\rm vir}\\sim 10^9 M_\\odot$ (e.g. Pe\\~narrubia et al. 2008a, Walker et al. 2009) and that in CDM the halo mass function is scale-free, all dark matter substructures orbiting in dwarfs are expected to fall below the star formation threshold and thus be dark. (iii) Dwarf galaxies are embedded in haloes with estimated dark matter densities that are several orders of magnitude higher than in the solar neighbourhood. In a scale-free cosmology this implies a higher number density of perturbers and thus a higher chance to detect their effects on a pre-existing binary population. (iv) The stellar populations of dSphs are generally dominated by old and intermediate-age components (e.g. Mateo 1998 and references therein), which enhances the cumulative effect of gravitational perturbations.  Unfortunaltey, although binaries with intermediate ($\\simless 100$ AU) separations have been studied in stellar clusters with some detail in the past (e.g. White \\& Ghez 2001; Patience et al. 2002), little is known about the number and distribution of binary stars in dwarf spheroidal galaxies (e.g. Wyse et al. 2002 estimate that 5\\% of the stars in the Ursa Minor dwarf correspond to unresolved binaries from the number of blue stragglers detected in this galaxy).  This paper is aimed at studying the disruption of wide ($\\simgreat 100$ AU) binaries as a result of interactions with dark substructures orbiting in dwarf galaxies, as well as exploring whether the effects can be detected with our present observational capabilities. The arrangement of the paper is as follows: \\S2 outlines the analytical analysis, \\S3 tests its consistency through N-body simulations, \\S4 applies the result to the Milky Way dwarf galaxy population, \\S5 estimates the feasibility of detecting wide binaries in dwarfs and \\S6 provides a summary of the results. ",
        "conclusions": "The present cosmological paradigm, $\\Lambda$CDM, predicts the presence of dark matter substructures devoided of baryonic material (the so-called ``dark galaxies'') in all galactic haloes on all mass scales. We have considered in this paper the possibility of using wide stellar binaries to uncover the existence of these objects in dwarf spheroidal galaxies (dSphs), whose estimated dark matter densities surpass those of any known galaxy in the Universe.  Our analytical estimates suggest that a large fraction of wide binaries can be wiped out due to tidal encounters with dark substructures. These interactions occur in a ''catastrophic'' regime, wherein the disruption rate is proportional to the density of perturbers averaged over the life-time of the binary population $\\langle \\rho_p \\rangle$.  We find that in this regime the perturbed binary distribution has a separation function that scales as $dN/da\\propto a^{-2.1}$ at $a\\simgreat a_{\\rm max}$, where $a_{\\rm max}\\propto \\langle \\rho_p \\rangle^{-2/3}$. Our estimates show that for the range of dynamical masses and Galactocentric distances spanned by Milky Way dSphs, {\\it all} dwarfs are expected to have a truncation in the separation function at $a_{\\rm max}\\simless 0.1 $pc within the CDM context.  We examine in which dSphs observations of large separation binaries are feasible with existing instruments, finding that the ACS camera on board of the Hubble telescope would be able to test the predictions enclosed in this paper in dSphs that locate at heliocentric distances $\\simless 100$ kpc, even if the binary fraction amounts only 10\\% of the stellar population.  Altogether these considerations may pose a strong test through which $\\Lambda$CDM shall soon have to pass."
    },
    "1005/1005.0752_arXiv.txt": {
        "abstract": "{We construct a model for delayed electroweak symmetry breaking that takes place in a cold Universe with $T\\ll 100\\,$GeV and which proceeds by a fast quench rather than by a conventional, slow, phase transition. This is achieved by coupling the Standard Model Higgs to an additional scalar field. We show that the quench transition can be made fast enough for successful Cold Electroweak Baryogenesis, while leaving known electroweak physics unchanged. } ",
        "introduction": "Introduction}  It is well known that electroweak baryogenesis within the Standard Model is subject to a double kill: the electroweak phase transition is a cross-over \\cite{Kajantie:1996mn} so that the requirement for non-equilibrium is never fulfilled; and at the electroweak phase transition temperatures, the CP-violation in the fermion mass matrix is too small by many orders of magnitude (see \\cite{Rubakov:1996vz} for a review). However, it has also been known for some time that the observed baryon asymmetry could be explained if the electroweak transition is cold rather than hot \\cite{GarciaBellido:1999sv,Krauss:1999ng,Copeland:2001qw,Tranberg:2003gi} with $T\\ll 100\\,$GeV, and if the transition itself is a fast quench rather than adiabatic \\cite{Tranberg:2006ip,Tranberg:2006dg}. Moreover, recent analytical work \\cite{Hernandez:2008db} and numerical simulations have indicated that in such a cold environment, Standard Model CP-violation is much larger than at electroweak temperatures and is indeed strong enough to account for the baryon asymmetry \\cite{Tranberg:2009de}. In Cold Electroweak Baryogenesis, most of the baryon asymmetry is produced at the initial quench when the Higgs field is rapidly falling down the slope of the potential. Baryon production essentially stops after the first oscillation, after which the coherent Higgs field starts decaying, thereby reheating the Universe. The reheating temperature should be low enough that sphaleron diffusion does not subsequently wipe out the baryon number.  In the early Universe the rate of any global change is necessarily related to the Hubble rate, which at the electroweak scale is very small. Hence, in the absence of a first order phase transition, the realization of a fast quench does not appear to be possible within the Standard Model. But if the Higgs field were coupled to some beyond-the-Standard Model fields, the situation might change. Indeed, the Higgs could be just one of many scalar fields such as in the low energy limit of string theory.  The inflaton field could also be an example of such ``moduli'', and a low temperature state and a fast quench could both follow from single field (\"inverted hybrid\"-type) low-scale inflation \\cite{Knox:1992iy,German:2001tz,Copeland:2001qw,vanTent:2004rc}. However, this is at the expense of allowing a rather general inflaton potential (up to the sixth power in the field), as well as some fine tuning of parameters. Presumably this tuning can be alleviated by considering more complicated models (see, for instance \\cite{Ross:2010fg}).  Alternatively, as in the present paper, one can assume that inflation is decoupled from the electroweak phase transition and takes place at some high energy scale. We then study the generic conditions for the cold, quenched electroweak phase transition in the presence of an extra scalar field coupled to the Standard Model Higgs. In particular, we focus on the possibility of triggering the quench through the cosmological expansion, via thermal corrections to the effective potential. We will demonstrate both that, under certain conditions, the phase transition can be delayed until $T\\simeq {\\cal O}(1)\\,$GeV, and that the transition can be fast (in a sense to be specified below).  We will describe briefly the problem of the fast quench in the next section. In section \\ref{sec:triggerhiggs} we couple the Higgs field to another scalar field and show that a fast quench can be achieved. In section \\ref{sec:triggersigma} we show how to trigger a quench of the $\\sigma$ through finite temperature corrections to the mass and the expansion of the Universe. We then briefly consider effects such as defect formation, bubble nucleation and loop corrections in section \\ref{sec:checks}. In section \\ref{sec:electroweak} we introduce the criteria to be fulfilled in order that the Standard Model physics remains unchanged and section \\ref{sec:theconclusion} contains our conclusions. ",
        "conclusions": "Conclusion}  Using a simple implementation of the Higgs-$\\sigma$ potential, we have argued that the expansion of the Universe can be responsible for triggering a fast electroweak symmetry breaking transition. At the same time, this transition can be delayed until the temperature of the Universe is far below the electroweak scale.  The model can be summarised in the potential \\be V(\\sigma,\\phi)=V_0-\\frac{m_\\sigma^2}{2}\\sigma^2-\\frac{\\lambda_4}{4}\\sigma^4+\\frac{\\lambda_6}{6}\\sigma^6-\\frac{g^2}{2}\\sigma^2\\phi^2+\\frac{m_\\phi^2}{2}\\phi^2+\\frac{\\lambda_\\phi}{4}\\phi^4+\\frac{\\kappa}{2}\\sigma^2\\xi^2+\\frac{m_\\xi^2}{2}\\xi^2,\\nonumber\\\\ \\ee and we split the evolution into 3 stages: \\begin{itemize} \\item 1) The post inflationary Universe in equilibrium cools under Hubble expansion until the $\\sigma$ is driven to a tachyonic transition. \\item 2) When $\\sigma_{\\rm ave}=\\langle\\sigma^2\\rangle\\simeq m_\\sigma^2/\\lambda_4$ Hubble expansion can ignored and $\\sigma$ rolls down its potential, triggering a fast tachyonic quench of the Standard Model Higgs, $\\phi$. \\item 3) $\\phi$ and $\\sigma$ settle in the global minimum of the potential, for which the physics is constrained by electroweak phenomenology. \\end{itemize}  \\EPSFIGURE[t]{./Pictures/final2.eps,width=9cm,clip}{The region in parameter space consistent with all three constraints presented in Figs.~\\ref{fig:vq}, \\ref{fig:m2} and \\ref{fig:gamma}. The dot is our example point $v_\\sigma=3000\\,$GeV, $m_\\phi=150\\,$GeV. \\label{fig:final} }  The cost is the addition of an extra scalar field coupling to some other matter sector ``$\\xi$'', and with self-couplings $\\sigma^4$ and $\\sigma^{6}$. This is partly a result of insisting that electroweak physics is unaltered, but it also comes from restricting ourselves to a minimal model. More complicated $\\sigma$-potentials can be envisaged, and for instance replacing ($\\sigma^4$, $\\sigma^6$) by ($\\sigma^5$, $\\sigma^6$) or ($\\sigma^6$, $\\sigma^8$) works as well. The input values used here for the Higgs mass ($m_H=160\\,$GeV) can be relaxed within the range still allowed by experiment, say $m_H\\in[115:200]\\,$GeV. Similarly, the height of the potential $V_0$ can be relaxed to $\\simeq 200^4\\,$GeV, while still avoiding finite temperature electroweak symmetry restoration after reheating.  The model is very similar to the low-scale inflation model proposed in \\cite{Copeland:2001qw,vanTent:2004rc}, although since now $\\sigma$ is not required to be responsible for inflation and the CMB fluctuations, a simpler potential and weaker tuning is allowed. Still, as in \\cite{vanTent:2004rc} there is $\\phi$-$\\sigma$ coupling and mixing, providing an observational signature of this model. However, in contrast to \\cite{vanTent:2004rc}, one mass eigenvalue ($m_-$) is smaller than $m_H$($\\simeq m_+$), and the Higgs can in principle decay into $\\sigma$ particles. This is however suppressed by a small $\\phi-\\sigma$ coupling.  We conclude that although the expansion of the Universe is very slow at the electroweak scale, it can be amplified in a rather straightforward way to provide a fast electroweak transition. This transition can be delayed indefinitely (essentially until $T\\simeq m_\\sigma$), and so also until the Universe is cold enough that the computation \\cite{Hernandez:2008db} of the effective Standard Model CP-violation becomes reliable. And so, although far from being the final word on the matter, we consider the scenario presented here an interesting option to realise a cold, fast electroweak transition."
    },
    "1005/1005.3685_arXiv.txt": {
        "abstract": "We present theoretical calculations for the differential distribution of stellar orbital eccentricity in a galaxy halo, assuming that the stars constitute a spherical, collisionless system in dynamical equilibrium with a dark matter halo.  In order to define the eccentricity $e$ of a halo star for given energy $E$ and angular momentum $L$, we adopt two types of gravitational potential, such as an isochrone potential and a Navarro-Frenk-White potential, that could form two ends covering in-between any realistic potential of dark matter halo. Based on a distribution function of the form $f(E,L)$ that allows constant anisotropy in velocity dispersions characterized by a parameter $\\beta$, we find that the eccentricity distribution is a monotonically increasing function of $e$ for the case of highly radially anisotropic velocity dispersions ($\\beta \\gtrsim 0.6$), while showing a hump-like shape for the cases from radial through tangential velocity anisotropy ($\\beta \\lesssim 0.6$). We also find that when the velocity anisotropy agrees with that observed for the Milky Way halo stars ($\\beta \\simeq 0.5-0.7$), a nearly linear eccentricity distribution of $N(e) \\propto e$ results at $e \\lesssim 0.7$, largely independent of the potential adopted. Our theoretical eccentricity distribution would be a vital tool of examining how far out in the halo the dynamical equilibrium has been achieved, through comparison with kinematics of halo stars sampled at greater distances. Given that large surveys of the SEGUE and Gaia projects would be in progress, we discuss how our results would serve as a new guide in exploring the formation and evolution of the Milky Way halo. ",
        "introduction": "\\label{intro}  Studies of large-scale structures in the universe and fluctuations in the cosmic microwave background strongly favor a $\\Lambda$-cold dark matter ($\\Lambda$CDM) cosmology (e.g., Cole et al. 2005; Dunkley et al. 2009). The formation of structures in this cosmology is a process of hierarchical clustering, in the sense that numerous CDM lumps cluster gravitationally and merge together to form larger structures (White \\& Rees 1978; Blumenthal et al. 1984).  Dark halos of galaxy systems are similarly formed via clustering of subhalos as a result of CDM agglomerations that reach the maximum expansion then turn around to collapse in the background expanding medium, but a detailed process leading to the halo formation from primordial density fluctuations is highly nonlinear and is not as simple as the formation of larger structures in the universe (e.g., for review see Ostriker 1993 and Bertschinger 1998).  High-resolution $\\Lambda$CDM simulations for the halo formation generically show that mergers and collisions of subhalos induce the overall collapse and virialize the inner region of host halo, while surviving subhalos orbit as separate entities within the inner virialized region of halo (e.g., Moore et al. 1999; Ghigna et al. 2000; Helmi, White \\& Springel 2003; Valluri et al. 2007). A majority of stars formed through this build-up of halo are expected to have also experienced the redistribution of energy and momentum that drives the phase mixing or violent relaxation towards the dynamical equilibrium (Lynden-Bell 1967). This leads to an idea that a stellar halo, which can be regarded as a collisionless system, holds the dynamical information just after the last violent relaxation in forming the halo.  We then take an approach to find out the relics of the formation of the Milky Way halo from the kinematics of halo stars.  Among many of their kinematic properties available at present and in the near future, the differential distribution of stellar orbital eccentricity $N(e)$ seems to be of special importance. The orbital eccentricity of a star is a quasi-adiabatic invariant (Eggen, Lynden-Bell \\& Sandage 1962; Lynden-Bell 1963) and is unaffected by the small and slow variation of the gravitational potential that might have occurred after the major formation of halo stars.  It is therefore most likely that the shape of $N(e)$ has been conserved until present.  With this consideration, comparing the observed shape of $N(e)$ for halo stars with the theoretical one for the halo in dynamical equilibrium, we could explore how far out in the halo the dynamical equilibrium was achieved. Consequently, $N(e)$ serves as a new test of halo formation scenario in a $\\Lambda$CDM cosmology.   As a useful way to derive $N(e)$ theoretically, we consider the orbit of halo stars in assumed gravitational potentials of the halo. In section 2, we present our formulation to calculate $N(e)$ under some plausible assumptions for the halo, and apply it to two extreme gravitational potentials of academic interest. The results for realistic cases are shown for the isochrone potential and for the Navarro-Frenk-White (NFW) potential in section 3.  We summarize the results and discuss the prospects of investigating the formation and evolution of the Milky Way halo in section 4.   \\section[]{Formulation} \\label{method} We assume that the halo stars constitute a spherical, collisionless system in dynamical equilibrium with a dark halo.  Since the dark matter is known to dominate the total mass of the galaxy system, the motion of halo stars is governed by the gravitational potential of dark halo.   \\subsection{Stellar orbital eccentricity in a model halo} \\label{e}   When a spherical halo potential $V(r)$ is given with respect to the galaxy center, the energy $E$ and the angular momentum $L$ of a star at the position ${\\bm r}$ with the velocity ${\\bm v}$ are written respectively as \\begin{equation} E=\\frac{1}{2} {\\bm v}^2 + V(r), \\hspace{2ex} {\\rm and} \\hspace{2ex} L = |{\\bm L}| = \\left| {\\bm r}\\times {\\bm v} \\right| , \\end{equation} where $r=|{\\bm r}|$. The orbital eccentricity of a star is practically defined as \\begin{equation} \\label{e-def} e \\equiv \\frac{ r_{\\rm apo} - r_{\\rm peri} }{ r_{\\rm apo} + r_{\\rm peri} } , \\end{equation} where $r_{\\rm apo}$ and $r_{\\rm peri}$ are the apo- and peri-centric distances, respectively, and are given by two real solutions ($r_{\\rm apo}>r_{\\rm peri}$) of the following equation: \\begin{equation} \\label{eq_r1r2} E = V(r) + \\frac{L^2}{2 r^2} \\equiv V_{\\rm{eff}} (L;r) . \\end{equation} It is evident from equations (\\ref{e-def}) and (\\ref{eq_r1r2}) that a pair of $(E,L)$ has a one-to-one correspondence to $(E,e)$, but there is a region of $(E,L)$ in which two real solutions are not allowed and thus, except for the case of circular orbits, the eccentricity cannot be defined. Since such unbound orbits do not form a steady population of stellar halo, we neglect them and exclusively consider stars with bound orbits.  Constraints on $E$ and $L$ that allow bound orbits are presented in Appendix A.   \\subsection{Differential distribution of stellar orbital eccentricity} \\label{d}  Let $f({\\bm r},{\\bm v})$ be the distribution function of halo stars, then the number of halo stars in a phase space volume $d^3 {\\bm r} d^3 {\\bm v}$ centered at $({\\bm r}, {\\bm v})$ is given by $f({\\bm r}, {\\bm v}) d^3{\\bm r}d^3{\\bm v}$. According to the strong Jeans theorem, the distribution function should be expressed in terms of isolating integrals only (Lynden-Bell 1960, 1962).  For a spherical system that is invariant under rotation, it takes a form of either $f(E)$ or $f(E,L)$, depending on whether the stellar velocity dispersion is isotropic or anisotropic, respectively.  The velocity dispersion observed for halo stars is radially anisotropic (e.g., Yoshii \\& Saio 1979; Gilmore, Wyse, \\& Kuijken 1989).  Furthermore, recent observations for halo stars within the distance of $10 \\hspace{1ex} {\\rm kpc}$ away from us show that the shape of velocity ellipsoid is constant and its principal axes are well aligned with the spherical coordinates (Carollo et al. 2007; Bond et al. 2009). If we extrapolate this fact to a whole system, one simple form of the distribution function is \\begin{equation}\\label{f_beta} f(E,L) = \\begin{cases} g(E) L^{-2 \\beta} , & \\text{if $(E,L)$ is `allowed'} \\\\ 0 ,                  & \\text{\\rm otherwise,} \\end{cases} \\end{equation} where $g(E)$ is a function of $E$ (e.g., Binney \\& Tremine 2008). Here, $\\beta$ is a constant value of velocity anisotropy parameter defined as \\begin{equation} \\beta \\equiv 1 - \\frac{\\sigma_{\\rm t}^2}{\\sigma_r^2}, \\end{equation} where $\\sigma_{r}$ is the radial velocity dispersion and $\\sigma_{t}$ is the tangential velocity dispersion projected onto the spherical $\\theta -\\phi$ surface. Although $\\beta$ is about $0.5-0.7$ observationally (e.g., Bond et al. 2009; Smith et al. 2009; Carollo et al. 2010), we will use it as a constant parameter below.  By changing the variables and integrating over the spherical coordinates, the number of stars in $d^3 {\\bm r} d^3 {\\bm v}$ reduces to \\begin{equation} \\label{with_L} N(E,L)dEdL^2 = 4\\pi^2 g(E) L^{-2\\beta} T_r dE dL^2 , \\end{equation} with the radial period of stellar orbit given by \\begin{equation}\\label{Tr} T_r (E,L) \\equiv \\oint \\frac{dr}{v_r} = 2 \\int_{r_{\\rm peri}}^{r_{\\rm{apo}}} \\frac{dr}{\\sqrt{2 \\left[ E-V_{\\rm{eff}}(L;r) \\right]}}. \\end{equation} Since $L^2$ is a function of $E$ and $e$, we here introduce the $E$-dependent differential eccentricity distribution as \\begin{equation} \\label{general-n} n_{\\beta}(E,e) = 4 \\pi^2 L^{-2 \\beta} \\cdot T_r \\cdot \\left|\\left(\\frac{\\partial L^2}{\\partial e}\\right)_{E}\\right| . \\end{equation} We then express the differential eccentricity distribution as \\begin{equation} \\label{general-N} N_{\\beta}(e) = \\int_{{\\rm allowed} \\; E} g(E) n_{\\beta}(E,e) dE . \\end{equation} It is apparent from this equation that $N_{\\beta}(e)$ is a weighted sum of $n_{\\beta}(E,e)$ with a weight function of $g(E)$.  Thus, once the gravitational potential $V(r)$ and the velocity anisotropy parameter $\\beta$ are specified, we can formally obtain $n_{\\beta}(E,e)$, and also $N_{\\beta}(e)$ after integrating $n_{\\beta}(E,e)$ over $E$ with its appropriate weight.   \\subsection{Extreme cases of mass distribution} \\label{extreme}  In this subsection, mostly for pedagogical purpose, we consider two extreme cases of mass distribution such as the point mass at the center and the homogeneous distribution in the truncated sphere.  These cases allow analytic expression of $n_{\\beta}(E,e)$, and because it is separable in $E$ and $e$, $N_{\\beta}(e)$ can also be obtained except for its normalization.  Therefore, these cases are helpful to understand the results for any more realistic cases.   \\subsubsection{Central point mass} \\label{central} The gravitational potential arising from the central point mass is Keplerian and is given by \\begin{equation} V(r) =  - \\frac{GM}{r} , \\end{equation} where $M$ is the total mass of dark halo and $G$ is the gravitational constant. For bound orbits with $E < 0$, there are two real and positive solutions for equation (\\ref{eq_r1r2}), or equivalently, \\begin{equation} (-2E)r^2 - 2GMr + L^2 = 0 . \\end{equation} The orbital eccentricity is expressed in terms of $(E,L)$ as \\begin{equation}\\label{1-e} e =  \\sqrt{1 - \\frac{(-2E)L^2}{{(GM)}^2}} , \\end{equation} and the other relevant quantities are neatly expressed as \\begin{equation} T_r = \\frac{2 \\pi GM}{ {(-2E)}^{\\frac{3}{2}} }, \\hspace{2ex} {\\rm and} \\hspace{2ex} L^2 = \\frac{{(GM)}^2}{-2E} (1-e^2). \\end{equation} Substitution of these quantities in equations (\\ref{general-n}) and (\\ref{general-N}) gives the $E$-dependent differential eccentricity distribution \\begin{equation} n_{\\beta} (E,e) = 16 \\pi^3 {(GM)}^{3-2\\beta} {(-2 E)}^{\\beta - \\frac{5}{2} } \\frac{e}{{(1-e^2)}^{\\beta}} , \\end{equation} and the differential eccentricity distribution \\begin{equation} N_{\\beta} (e) = 16 \\pi^3 {(GM)}^{3-2\\beta} \\left[\\int g(E){(-2 E)}^{\\beta - \\frac{5}{2} }dE\\right] \\frac{e}{{(1-e^2)}^{\\beta}} . \\end{equation} Since $N_{\\beta}(e)\\propto n_{\\beta} (E,e)$, we normalize $N_{\\beta}(e)$ such that $\\int_{0}^{1} N_{\\beta}(e)de=1$, and write \\begin{equation} {\\rm normalized}\\;\\; N_{\\beta}(e)=2(1-\\beta) \\frac{e}{{(1-e^2)}^{\\beta}}, \\hspace{2ex} (\\beta \\neq 1) . \\end{equation} The results of $N_{\\beta}(e)$ for several values of $\\beta$ are shown on the left panel of Figure \\ref{Fig Kepler+IHO-beta}. For the case of $\\beta=0$ (isotropic velocity dispersion), $N_{\\beta}(e)$ is exactly proportional to $e$ (Binney \\& Tremaine 2008) and we call it the linear eccentricity distribution.  For $0<\\beta<1$ (radially anisotropic velocity dispersion), $N_{\\beta}(e)$ is a rapidly increasing function of $e$ with a peak always at $e=1$.  On the other hand, for $\\beta<0$ (tangentially anisotropic velocity dispersion), $N_{\\beta}(e)$ shows a hump-like $e$-distribution around a single peak at $e=(1-2\\beta)^{-1/2}$.  We should notice that the linear trend of $N_{\\beta}\\propto e$ prevails in a range of $0<e<0.3$ regardless of $\\beta$, while the behavior of $N_{\\beta}(e)$ is very sensitive to $\\beta$ in a range of $0.6<e<1$ and the difference there clearly shows up.  \\begin{figure*} \\begin{center} \\includegraphics[width=\\columnwidth]{Fig1a.eps} \\includegraphics[width=\\columnwidth]{Fig1b.eps}\\\\ \\caption{Differential distribution of stellar orbital eccentricity $N_{\\beta}(e)$ in two extreme cases of mass distribution, such as the point mass model on the left panel and the homogeneous model on the right panel.  The results are shown by lines for several values of velocity anisotropy parameter $\\beta$.  If $N_{\\beta}(e)$ near $e=1$ sensitively changes at some particular value of $\\beta$, the results for $\\beta\\pm 0.05$ are additionally shown by dotted lines for the purpose of illustrating its sensitivity. Note that $N_{\\beta}(e)$ is normalized such that $\\int_{0}^{1} N_{\\beta}(e)de=1$. } \\label{Fig Kepler+IHO-beta} \\end{center} \\end{figure*}   \\subsubsection{Truncated homogeneous sphere} \\label{homo}  A homogeneous density distribution within truncated sphere is expressed as \\begin{equation} \\rho(r) = \\begin{cases} \\frac{3M}{4 \\pi r_{\\rm t}^3} , & \\text{if $r < r_{\\rm t}$} \\\\ 0 ,                      & \\text{otherwise,} \\end{cases} \\end{equation} where $M$ is the total mass of dark halo and $r_{\\rm t}$ is the truncation radius. The gravitational potential arising from this density distribution is given by \\begin{equation} V(r) = \\begin{cases} - \\frac{3GM}{2 r_{\\rm t}} + \\frac{GM}{2 r_{\\rm t}} {\\left( \\frac{r}{r_{\\rm t}} \\right) }^2 , & \\text{if $r < r_{\\rm t}$} \\\\ - \\frac{GM}{r} , & \\text{otherwise.} \\end{cases} \\end{equation} We consider only stars with $E<E_{\\rm t}\\equiv-GM/r_{\\rm t}$, which guarantees the stars to be confined inside the truncated radius $r_{\\rm t}$. Thus, bound orbits within the truncated sphere are allowed if $E_{\\rm min}<E<E_{\\rm t}$ where we note $E_{\\rm min} \\equiv (3/2)E_{\\rm t}$. In this limited range of $E$, there are two real and positive solutions for equation (\\ref{eq_r1r2}), or equivalently, \\begin{equation} GMr_{\\rm t} {\\left( \\frac{r}{r_{\\rm t}} \\right)}^4 - 2 \\left( E - E_{\\rm min} \\right) r_{\\rm t}^2 {\\left( \\frac{r}{r_{\\rm t}} \\right)}^2+ L^2 = 0 , \\end{equation} if and only if \\begin{equation} 0 < D < 1, \\end{equation} where \\begin{equation} D = \\frac{GML^2}{r_{\\rm t}^3  {\\left( E -E_{\\rm min} \\right)}^2 } . \\end{equation} The orbital eccentricity is expressed in terms of $D$ as \\begin{equation}\\label{2-Ie} e = \\sqrt{\\frac{1-\\sqrt{D}}{1+\\sqrt{D}}} , \\end{equation} and the other relevant quantities are expressed in terms of $(E,L)$ as \\begin{equation} T_r = \\pi \\sqrt{ \\frac{r_{\\rm t}^3}{GM} }, \\hspace{2ex} {\\rm and} \\hspace{2ex} L^2 = \\frac{r_{\\rm t}^3}{GM} {\\left( E - E_{\\rm min} \\right)}^2 { \\left(\\frac{1-e^2}{1+e^2}\\right) }^2 . \\end{equation} Consequently, we obtain \\begin{equation} \\label{nIHO} n_{\\beta}(E,e) = 32 \\pi^3 {\\left( \\frac{r_{\\rm t}^3}{GM} \\right)}^{\\frac{3}{2} - \\beta} {\\left( E - E_{\\rm min} \\right)}^{2-2\\beta} \\frac{e {(1-e^2)}^{1-2\\beta}}{{(1+e^2)}^{3-2\\beta}} , \\end{equation} and \\begin{multline} N_{\\beta}(e) = 32 \\pi^3 {\\left( \\frac{r_{\\rm t}^3}{GM} \\right)}^{\\frac{3}{2} - \\beta} \\left[ \\int g(E) {\\left( E - E_{\\rm min} \\right)}^{2-2\\beta}dE \\right] \\\\ \\times \\frac{e {(1-e^2)}^{1-2\\beta}}{{(1+e^2)}^{3-2\\beta}} . \\end{multline} As in the point mass model, since $N_{\\beta}(e)\\propto n_{\\beta} (E,e)$, we normalize $N_{\\beta}(e)$ such that $\\int_{0}^{1} N_{\\beta}(e)de=1$, and write \\begin{equation} {\\rm normalized}\\;\\; N_{\\beta}(e)=8(1-\\beta) \\frac{e {(1-e^2)}^{1-2\\beta}}{{(1+e^2)}^{3-2\\beta}} \\hspace{2ex}, (\\beta \\neq 1) . \\end{equation}   The results of $N_{\\beta}(e)$ for several values of $\\beta$ are shown on the right panel of Figure \\ref{Fig Kepler+IHO-beta}. For $\\beta<0.5$, $N_{\\beta}(e)$ shows a hump-like $e$-distribution with a single peak at \\begin{equation} e_{\\rm peak}= \\sqrt{\\frac{ 4(1-\\beta) - \\sqrt{13 - 32\\beta + 16\\beta^2} }{3} } . \\end{equation} For $0.5<\\beta<1-\\sqrt{3}/4$, however, $N_{\\beta}(e)$ has two local maxima such as a broad peak at $e=e_{\\rm peak}$ and a sharp peak at $e=1$.  Overall behavior monotonically increases with $e$ in a range of $0<e<e_{\\rm peak}$, and is kept more or less flat in the range of $e_{\\rm peak}<e<1$. For $1-\\sqrt{3}/4<\\beta<1$, $N_{\\beta}(e)$ is a rapidly increasing function of $e$.  For a given value of $\\beta$, $N_{\\beta}(e)$ is more weighted at smaller $e$ in the homogeneous model, when compared with the point mass model.  In particular, for $\\beta=0$ (isotropic velocity dispersion), $N_{\\beta}(e)$ shows a broad hump-like $e$-distribution around a peak at $e_{\\rm peak}=0.36$ in the homogeneous model, while showing an exactly linear $e$-distribution in the point mass model.  This sensitivity, though between two extreme cases, could be used to discriminate the likely mass distribution in more realistic cases to be considered in section 3.    \\subsection{Effect of central mass concentration} \\label{effect}  In the cases of central point mass and truncated homogeneous sphere, the shape of $n_{\\beta}(E,e)$ is the same as $N_{\\beta}(e)$, because $n_{\\beta}(E,e)$ is separable in $E$ and $e$ and thus the shape of $N_{\\beta}(e)$ is unaffected by $g(E)$ in equation (\\ref{f_beta}).  This property generally holds when the density distribution in the truncated sphere is given by $\\rho (r) \\propto 1/r^{\\gamma}$ (see Appendix B).  The homogeneous model in section \\ref{homo} corresponds to $\\gamma = 0$.  Using the cases of $\\gamma =1$ (linear potential model) and $\\gamma =2$ (singular isothermal model) that are intermediate between two extreme cases considered above, we can examine how $N_{\\beta}(e)$ depends on the central mass concentration. As shown in Figure \\ref{Fig various-models} for $\\beta = 0$, there is a clear trend that the $e$-distribution is peaked at larger $e$ as the halo mass is more centrally concentrated.  This trend is also true regardless of the value of $\\beta$ and is helpful in interpreting the results of more realistic models in the next section.   \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth]{Fig2.eps} \\caption{Differential distribution of stellar orbital eccentricity $N_{\\beta}(e)$ for four types of model potentials, when the stellar velocity dispersion is isotropic ($\\beta=0$).  Considered are the homogeneous model ($\\gamma =0$, section \\ref{homo}), the linear potential model ($\\gamma =1$, Appendix B.1), the singular isothermal model ($\\gamma =2$, Appendix B.2), and the point mass model (section \\ref{central}), in order of increasing the central mass concentration. There is a clear trend that the $e$-distribution is peaked at larger $e$ as the halo mass is more centrally concentrated. Note that the distribution is normalized such that $\\int_{0}^{1} N_{\\beta}(e)de=1$. } \\label{Fig various-models} \\end{center} \\end{figure} ",
        "conclusions": "\\label{discussion}  Hierarchical clustering scenarios of galaxy formation suggest that the major merger of at least several subhalos with comparable masses would occur at the last stage of galaxy formation.  This last major merger would cause the violent relaxation of halo stars and make them in dynamical equilibrium with a dark halo. Based on the assumptions that approximate such a status just after the last violent relaxation (section \\ref{intro}), we have presented theoretical predictions of $N(e)$ for halo stars. This predicted $N(e)$ should be observed for the Milky Way halo if it is an isolated system and the subsequent variation of the potential is quiescent enough to conserve the eccentricity of each star.  However, recent nearby observations suggest that at least some part of the Milky Way halo may have originated from accreted satellites, which possibly deviates the observed $N(e)$ from our predictions. For example, if infalling satellites break up and spread their member stars into the field, these stars would show peculiar eccentricity distribution which necessarily imprints the initial condition of the progenitor satellites. In addition, if such satellites locally disturb the halo potential, some {\\it in-situ} halo stars may have altered their orbits (e.g. Zolotov et al. 2009). With an invention of segregating  {\\it in-situ} halo stars from infalling stars, we might be able to well understand the nature of accretion and distortion of satellites.   Numerous authors subdivided halo stars into some `components' and examined the correlations between chemistry, age and kinematics of stars in each component. Carollo et al. (2010) obtained reliable eccentricities for $\\sim 10,000$ halo stars within $4\\;{\\rm kpc}$ of the sun and decomposed them into the inner and outer halo components having distinct  eccentricity distributions from each other. Since their sample is {\\it local} and is inherently biased in favour of stars that stay longer in the surveyed region, our formalism, which is designed to predict $N(e)$ of the whole stellar halo, has to be modified for the purpose of fair comparison with their data. Through proper incorporation of effects of such a bias, we can still predict $N(e)$ for a {\\it local} sample by fully taking into account a probability of finding each of halo stars in the surveyed region. This will be done in a separate paper in preparation. On the other hand, our formalism can directly apply to a {\\it global}, and therefore less biased, sample of halo stars with reliable orbital eccentricities, such as those from next generation surveys including the Gaia mission. In either case, the analytical approach in the present paper certainly forms a basis that serves as a useful tool for analysing the kinematics of the stellar halo.  Large, unbiased database of halo stars would enable us to test whether a given component is in dynamical equilibrium by comparing the observed and predicted shape of $N(e)$. Such comparison would hopefully discover some relaxed components, and their adiabatically conserved shape of $N(e)$ would carry some useful information of the physics of violent relaxation. Moreover, the spatial distribution of these relaxed components would enable us to see how far out in the halo the violent relaxation has exerted and how strongly it has affected the stellar halo.  If the information of last violent relaxation, yet to be known observationally, is gained in this way, more precise assessment to the early evolution of the Milky Way would be possible, and our understanding of its formation would greatly be advanced.  Our current calculations of $N(e)$ are certainly very simple and can be improved by using more realistic assumptions. For example, we can modify our analysis to allow axisymmetric potentials including a disk-like component as well as a bulge. Preliminary analysis has confirmed that inclusion of a disk-like component would cause no significant change in the linear trend of $N(e)$ described in section \\ref{result}, which will be discussed in a separate paper. Also, our choice of $f(E,L)$ having the form in equation (\\ref{f_beta}) has to be extended to allow the radial dependence of $\\beta(r)$. Further elaborate modeling of $N(e)$ with these theoretical improvements, when applied to future large survey of halo stars, would then provide a promising way of unraveling mysteries of the galaxy formation and  evolution in a paradigm of hierarchical clustering in the $\\Lambda$CDM cosmology."
    },
    "1005/1005.1649_arXiv.txt": {
        "abstract": "In this article, I briefly review the status of infrared effects which occur when using inflationary models to calculate initial conditions for a subsequent hot, dense plasma phase. Three types of divergence have been identified in the literature: secular, ``time-dependent'' logarithms, which grow with time spent outside the horizon; ``box-cutoff'' logarithms, which encode a dependence on the infrared cutoff when calculating in a finite-size box; and ``quantum'' logarithms, which depend on the ratio of a scale characterizing new physics to the scale of whatever process is under consideration, and whose interpretation is the same as conventional field theory. I review the calculations in which these divergences appear, and discuss the methods which have been developed to deal with them.  \\vspace{3mm} \\begin{flushleft} \\textsf{\\textbf{Keywords}: Inflation, Cosmological perturbation theory, Physics of the early universe, Quantum field theory in curved spacetime.} \\end{flushleft} ",
        "introduction": "\\label{sec:intro}  Looking out at the universe from the vantage point of Earth, we see a small fluctuation $\\delta T$ in the temperature of the cosmic microwave background radiation (CMB). Over the last decade, advances in observational astronomy have allowed us to commence a detailed study of $\\delta T$'s statistical properties. We aim to compare these properties with those of small fluctuations which, according to our present ideas, are expected to have existed in the very early universe.  These fluctuations evolve in a hot, dense plasma of tightly coupled baryons and photons. If we suppose the plasma era was preceded by an epoch of primordial inflation, then initial conditions for these fluctuations can be calculated from the parameters of the inflationary model. The variance predicted by this approach was computed in a now-classic series of early papers \\cite{Bardeen:1983qw,Guth:1982ec,Hawking:1982cz,Hawking:1982my, Mukhanov:1985rz,Sasaki:1986hm}, and later extended to the skewness \\cite{Allen:1987vq,Falk:1992sf,Gangui:1993tt,Maldacena:2002vr, Seery:2005gb,Chen:2006nt} and kurtosis \\cite{Seery:2006vu,Arroja:2008ga,Seery:2008ax,Chen:2009bc}. With sufficiently precise experimental data in hand, we can work backwards from observations of these statistical properties to determine what the initial conditions must have been. The result is a realistic prospect of constraining parameters in certain inflationary models.  This ambitious programme was begun in earnest following the arrival of high-quality CMB maps made by the \\emph{Wilkinson Microwave Anisotropy Probe} (WMAP) \\cite{Komatsu:2010fb}. Even more precise data are expected from the \\emph{Planck Surveyor} satellite. Once it had become clear that the statistical properties of $\\delta T$ carried important information about whatever physics was operative during inflation, theorists were soon tempted to refine their calculations in the hope that additional detail could be extracted. A similar programme in electroweak physics had provided important clues about the unknown details of electroweak symmetry breaking,% \\footnote{More details of this ``precision electroweak programme'' can be found in the review \\emph{Electroweak model and constraints on new physics,} prepared by the Particle Data Group \\cite{Amsler:2008zzb}.} and it was reasonable to explore the possibility of similar gains in cosmology. The observable statistical properties of the inflationary density perturbation are encoded in its $n$-point correlation functions, with the leading contribution to each function typically suppressed by $(n-1)$ powers of the small quantity $(H/\\Mp)^2$ \\cite{Jarnhus:2007ia}, where $H$ is the Hubble rate during inflation and $\\Mp = (8\\pi G)^{-1/2}$ is the Planck mass. Two types of refinement were possible: either to calculate to lowest order in $(H/\\Mp)^2$ at progressively larger $n$, or to calculate higher order corrections with $n$ fixed.  It had been known for a long time that perturbation theory beyond leading order in $(H/\\Mp)^2$ was complicated by troublesome infrared behaviour. Sasaki, Suzuki, Yamamoto \\& Yokoyama considered a $\\lambda \\phi^n$ theory with $n \\geq 3$, coupled non-minimally to gravity, and evaluated each $n$-point function perturbatively \\cite{Sasaki:1992ux,Suzuki:1992gi}. Sasaki {\\etal} noticed that the $n$-point functions evaluated in de Sitter space had a potential divergence in the far future, which they called ``superexpansionary.'' We will rephrase the analysis of Sasaki {\\etal} in modern notation and discuss the meaning of these divergences in \\S\\ref{sec:zoo-time} and \\S\\ref{sec:time}.  Later, Mukhanov, Abramo \\& Brandenberger studied a different sort of infrared effect---% the accumulation of long wavelength fluctuations in an expanding universe \\cite{Mukhanov:1996ak,Abramo:1997hu}. (See also Abramo \\& Woodard, Ref.~\\cite{Abramo:1998hi}.) This differed from the analysis of Sasaki {\\etal}, which had been based on a purely geometrical effect. Whether such divergences are physical is complicated by the question of gauge, initially studied by Unruh \\cite{Unruh:1998ic}. Later, Losic \\& Unruh were able to give a gauge-invariant argument that infrared terms generically become large in spacetimes close to de Sitter \\cite{Losic:2005vg,Losic:2006ht,Losic:2008ht,Unruh:2008zza}. We will discuss corrections of this type in \\S\\ref{sec:zoo-box} and \\S\\ref{sec:space}. Divergences of a third kind had long been studied by Prokopec, Tsamis, Woodard and their collaborators. These authors took the quantum nature of correlation functions seriously, and calculated loop corrections just as one would when studying scattering experiments. Unfortunately, the complexity of both the rules for calculating these loops and the computations themselves are somewhat greater than in Minkowski space. These calculations have been reported in a large literature. Together with quantum effects studied by others, including Boyanovsky, de Vega \\& Sanchez, a recent sample can be found in Refs.~\\cite{Prokopec:2002uw,Onemli:2002hr,Prokopec:2003bx,Onemli:2004mb, Brunier:2004sb, Boyanovsky:2005sh,Boyanovsky:2005px,Kahya:2005kj, Kahya:2006hc,Bilandzic:2007nb, Kahya:2006ui,Prokopec:2006ue}. We will consider effects of this type in \\S\\ref{sec:zoo-new}.  Perturbative calculations of correlation functions from inflation had been studied in the 1980s by Allen, Grinstein \\& Wise \\cite{Allen:1987vq} and later by Falk, Rangarajan \\& Srednicki \\cite{Falk:1992sf}. The sophistication of such calculations increased dramatically in the years following Maldacena's calculation of the three-point function in single-field slow-roll inflation \\cite{Maldacena:2002vr}. At least in part, this increase was a consequence of the more complicated models of inflation which had been developed during the 1990s and early 2000s, and which theorists hoped would yield a distinctive pattern of correlations \\cite{Seery:2005wm,Chen:2006nt,Huang:2006eha,Chen:2009bc}. These developments are reviewed pedagogically in a recent article by Chen \\cite{Chen:2010xk}.  Since 2002, significant theoretical effort has been expended in refining our understanding of both inflationary correlation functions themselves, and the models which can generate significantly non-gaussian statistics. The calculations reviewed in this paper were a consequence of the reinvigoration of interest in infrared issues which followed this effort. Logarithms exhibiting secular growth with time were discovered in calculations of $n$-point correlation functions with $n \\geq 3$ \\cite{Falk:1992sf,Zaldarriaga:2003my, Sloth:2006az,Sloth:2006nu,Seery:2008qj}. Other calculations, if taken at face value, gave divergent answers in the infinite-volume limit. Divergences could be avoided by carrying out the calculation in a finite box, but any such procedure left behind residual logarithms involving the box cutoff \\cite{Boubekeur:2005fj,Lyth:2005fi, Sloth:2006az,Sloth:2006nu,Byrnes:2007tm,Lyth:2007jh,Seery:2007wf}. Further logarithms were encountered while studying the influence of unknown physics at energy scales even larger than the inflationary scale \\cite{Weinberg:2005vy,Weinberg:2006ac,Chaicherdsakul:2006ui,Seery:2007we, Adshead:2008gk,Dimastrogiovanni:2008af,Seery:2008ms,Senatore:2009cf}.  Infrared effects are not a phenomenon unique to de Sitter calculations, but are present in many applications of quantum field theory. Their appearance typically signals the presence of a non-trivial background, which cannot be described by asymptotic \\emph{in}- and \\emph{out}-states containing a definite number of particles. In scattering calculations, infrared divergences allow putative fixed-particle states to develop a slowly varying field, made out of an accumulation of arbitrarily many soft particles radiated on approach to or recoil from the scattering event. This ``Bloch--Nordsieck'' or ``initial state radiation'' phenomenon is studied in many textbooks on quantum field theory \\cite{Weinberg:1995mt}. A similar effect can occur in the confined phase of quantum chromodynamics, where soft particles can be radiated by partons moving inside nuclei. These particles are trapped by the strong QCD force, and change the background colour field. This effect may dramatically alter the mix of partons observed by an impinging probe such as a photon, and cannot be neglected in an accurate comparison with experiment \\cite{Gribov:1972ri,Gribov:1972rt,Altarelli:1977zs,Kogut:1974ni, Dokshitzer:1977sg,Collins:1988wj}. The effective $W$ approximation is a related example, in which the same infrared effects control the $W^\\pm$ content of colliding nuclei. The results are very practical. For example, this method can be used to predict the rate of Higgs production from $W^+ W^-$ fusion \\cite{Gunion:1989we}.  Throughout this article, Planck's constant $\\hbar$ and the speed of light $c$ are set to unity. Many expressions are written in terms of the reduced Planck mass, $\\Mp^{-2} = 8 \\pi G$. For brevity, a scalar field which contributes to the energy--momentum tensor $T_{ab}$ is referred to as `active,' whereas a field making no contribution to $T_{ab}$ is referred to as a `spectator.' In a cosmological context, such scalars are isocurvature modes. ",
        "conclusions": "Infrared effects have attracted considerable interest, but at the time of writing it has not been possible to extract a definitive prediction of a novel effect from these calculations---that is, an effect which is not already implicit in calculations carried out locally around background fields set at a point on the scalar potential.  The best studied examples of infrared effects are the time-dependent secular logarithms introduced in \\S\\ref{sec:zoo-time}. Over short timescales these logarithms represent evolution of correlations, as the background fields slowly roll down their potentials. At later times quantum effects can become important, but their role remains somewhat unclear. The calculation of Burgess {\\etal}, supported by the arguments of Senatore \\& Zaldarriaga, suggests that in stable theories they will suppress correlations at late times.  Box-cutoff logarithms represent a different problem: the possibility that, after many e-folds of inflation, spatial gradients develop in the background scalar fields. If we only wish to use our theories of inflation to study observable correlations generated on approach to reheating, then this may not be a significant problem. On the other hand, if we wish to study the evolution of a large spatial volume undergoing inflation for a long period of time, we expose ourselves to the possibility of unknown new physics at large values of $H/\\Mp$. The gauge invariant arguments of Losic \\& Unruh appear to cast doubt on the idea that perturbation theory remains globally well defined for an indefinite period in de Sitter or nearly de Sitter spacetimes.  Except in special cases, such as single-field inflation, the study of correlations generated near exit from inflation is still ultraviolet-sensitive (in the sense of physics at large values of $H/\\Mp$) and infrared-sensitive (in the sense of physics on very large lengthscales), because we do not know in which reheating minimum inflation is likely to terminate. This need not be a problem if observation can accurately determine the properties of our local minimum. In that case we would be able to understand all the properties of our vacuum, but we would not be able to understand how to embed this minimum in the larger landscape of the inflationary potential. However, this problem is not new and has been with us since the early days of the inflationary paradigm.  \\ack  My understanding of the ideas discussed in this article is due to discussions with my colleagues and collaborators. I would like to thank Peter Adshead, Nicola Bartolo, Christian Byrnes, Xingang Chen, Emanuela Dimastrogiovanni, Richard Easther, Louis Leblond, James Lidsey, Eugene Lim, Bojan Losic, David Lyth, Karim Malik, Sami Nurmi, Dmitry Podolsky, Antonio Riotto, Sarah Shandera, Martin Sloth, Andrew Tolley, Filippo Vernizzi and Daniel Wesley for sharing many insights concerning inflation generally, and the treatment of de Sitter loop diagrams in particular. Clifford Burgess, Xingang Chen, Richard Holman, Louis Leblond, Sarah Shandera and Martin Sloth offered helpful criticisms on an early typescript of this review. I would like to thank Peter Adshead for a significant investment of time in verifying my recomputations of the inflaton self-loop, spectator-loop and graviton-loop diagrams discussed in \\S\\ref{sec:zoo-new}.  \\end{fmffile}"
    },
    "1005/1005.3827_arXiv.txt": {
        "abstract": "{ Orion is the nearest star-forming region that hosts a significant number of young and massive stars. The energy injected by these OB stars is thought to have created the Eridanus superbubble. Because of its proximity, Orion is a prime target for a detailed investigation of the interaction between massive stars and their environment. } { We study the massive star population of Orion and its feedback in terms of energy and mass, in order to compare the current knowledge of massive stars with kinematic and radioactive tracers in the surrounding interstellar medium (ISM). } { We assemble a census of the most massive stars in Orion, then use stellar isochrones to estimate their masses and ages, and use these results to establish the stellar content of Orion's individual OB associations. From this, our new population synthesis code is utilized to derive the history of the emission of UV radiation and kinetic energy of the material ejected by the massive stars, and also follow the ejection of the long-lived radioactive isotopes $^{26}$Al and $^{60}$Fe. In order to estimate the precision of our method, we compare and contrast three distinct representations of the massive stars. We compare the expected outputs with observations of  $^{26}$Al gamma-ray signal and the extent of the Eridanus cavity. } { We find an integrated kinetic energy emitted by the massive stars of 1.8$^{+1.5}_{-0.4}\\times10^{52}$ erg. This number is consistent with the energy thought to be required to create the Eridanus superbubble. We also find good agreement between our model and the observed $^{26}$Al signal, estimating a mass of $5.8^{+2.7}_{-2.5}\\times10^{-4}$ $M_{\\odot}$ of $^{26}$Al in the Orion region. } { Our population synthesis approach is demonstrated for the Orion region to reproduce three different kinds of observable outputs from massive stars in a consistent manner: Kinetic energy as manifested in ISM excavation, ionization as manifested in free-free emission, and nucleosynthesis ejecta as manifested in radioactivity gamma-rays. The good match between our model and the observables does not argue for considerable modifications of mass loss. If clumping effects turn out to be strong, other processes would need to be identified to compensate for their impact on massive-star outputs. Our population synthesis analysis jointly treats kinematic output and the return of radioactive isotopes, which proves a powerful extension of the methodology that constrains feedback from massive stars. } ",
        "introduction": "At a distance of only some 400 pc, the Orion region is sufficiently close that it enables us to study its stellar population and interstellar-gas morphology in detail. The massive star ($M$ $>$ 8$M_{\\odot}$) population  is dominated by the Orion OB1 association, which includes four subgroups labelled a-d \\citep{Brown1994}. Their ages have been estimated to range between 1 and 12 Myr. OB1 is located on the near side of the densest part of the Orion Molecular clouds \\citep{Maddalena1986}, facing the Eridanus cavity which extends from these molecular clouds towards the Sun. H$\\alpha$ features, which coincide with a hole in the HI distribution \\citep{Heiles1976}, together with X-ray emission near HI features \\citep{Burrows1993} outline this large interstellar cavity, and provide evidence for the interactions between the hot gas in the Eridanus cavity and the neutral surrounding interstellar medium \\citep{Bally1991}. Many different names are used for the various substructures of Orion. We use the same convention as \\citet{Bally2008} in his recent review.   Feedback from massive stars plays a crucial role in the formation of stars, as it is shaping the ISM and its subsequent star formation activity. The main feedback originates from the ejection of matter from massive stars through their winds and supernova explosions, and from their intense emission at short wavelengths into the UV. This UV radiation creates large photoionized regions around the stars, and the kinetic energy associated with ejection of stellar matter pushes at the ISM, together creating large shells and cavities \\citep[e.g.][]{Heiles1976,vanderHucht1987,Leitherer1992,Maeder1994}. Kinetic energy output and UV radiation of massive stars were studied in \\citet{Voss-popsyn}, discussing the total emission from a population of stars with emphasis on the differences between various alternative stellar models. It was found that the kinetic energy from winds dominates over the supernova contribution, when integrated over the first 10 Myr after the stars were formed \\citep[see also][]{Leitherer1999}. This is due to the high wind velocities  % \\citep[taken from][]{Howarth,Lamers1995} and high mass-loss rates of the most massive stars \\citep{Castor1975} and \\citep{Vink2000}, even when modest wind clumping (with clumping factors of about 5) is accounted for \\citep{Repolust2004,Mokiem2007}. To understand the complex interplay between the massive stars and their local environment, which eventually result in the evolution of disks in galaxies, it is necessary to first create a census of the radiation, energy, and matter output of individual nearby star-forming regions like Orion, which can be validated in terms of observational constraints.  The radioactive isotope $^{26}$Al provides an interesting independent view on the interaction between young stars and the surrounding environment. It is traced by its $\\gamma$-ray decay line at 1808.63 keV, which can be observed with $\\gamma$-ray telescopes. With a mean lifetime of $\\sim$1 Myr $^{26}$Al is a long-term tracer of nucleosynthesis from populations of massive-star sources, as they eject it after synthesis in stellar cores and the supernova itself \\citep{Prantzos1996}. Typically, massive stars eject a few $10^{-5}M_{\\odot}$ of $^{26}$Al through their winds and supernovae (SN) \\citep[see e.g.][]{Limongi2006}. From $\\gamma$-ray observations, the total mass of $^{26}$Al in the Milky Way is estimated to be 2.8$\\pm$0.8$M_{\\odot}$ \\citep{Diehl2006}. Measurements of the $^{26}$Al emission from Orion by the COMPTEL instrument on NASAs Compton observatory generally confirms this scenario, with a $\\gamma$-ray intensity of $\\sim$ 7.5~10$^{-5}$~ph~cm$^{-2}$ s$^{-1}$ \\citep{Diehl2002}. The map of this $^{26}$Al $\\gamma$-ray emission, though not significant in its details, shows an interesting offset of the $^{26}$Al emission from the massive stars that are believed to the the source of the $^{26}$Al, and the emission appears rather extended \\citep{Diehl2002}. This suggests that the radioactive ejecta stream  into the nearby Eridanus cavity from their stellar association sources. A similar tracer would be the isotope $^{60}$Fe, observed in the Galaxy globally by its 1173 keV and 1333 keV decay lines \\citep{Smith2005,Harris2005,Wang2007}. This isotope is presumably created in neutron capture reactions in late shell burning stages of such massive stars, and is also emitted in the supernova explosions \\citep{Limongi2006}. $^{60}$Fe has a mean lifetime of $\\sim$3.6 Myr \\citep{Rugel2009}. It has not been seen from the Orion region, which is however not surprising, as its $\\gamma$-ray intensity has been found to be $\\sim$15\\% of the $^{26}$Al $\\gamma$-ray intensity only.  In this paper we analyze the energy and radioactive-isotope output from the entire stellar content of the Orion region. We compare the results with observational constraints, such as the measured strength of the 1808.63 keV line from $^{26}$Al decay and the size of the Eridanus superbubble. ",
        "conclusions": "% We analyzed the population of massive stars in the nearby star-forming Orion region, including the four OB1 subgroups (a-d) and the $\\lambda$ Ori group. We analyzed the stellar contents of the individual groups, providing updated lists of the stars more massive than 8 $M_{\\odot}$. Ages of the individual groups were constrained based on comparison between the updated properties of the most massive stars and stellar isochrones. Based on these results, we performed a study of the ejection of kinetic energy and radioactive elements from the young massive stars in Orion. We showed that the current state of the region only depends modestly on the properties of the model, such as the star formation history and the stellar evolution models. Main uncertainties are due to the unknown population of very massive stars that exploded over the past 10 Myr.  The population synthesis results were compared to the energy needed to form the Eridanus superbubble, the emission of hydrogen ionizing photons, and the intensity of the 1.809 MeV line from the decay of $^{26}$Al, showing good agreement between our model estimates and the observations. The $^{26}$Al observations provide a valuable tracer of the population of (now not any more observable) stars and thus of the cumulative action of massive star groups, and of the kinematics of the outflows from the massive stars.  Our current understanding of stellar evolution and supernova models is far from complete. Different models often rely on similar assumptions. Showing consistency between models and observations is important, as it supports confidence that the most important effects are accounted for in models. We have employed different models for characterizing the Orion region's stellar population, and for the stellar-evolution inputs to population synthesis. The results show that the observed properties of the Orion region are consistent with these models. Differences among models are smaller than the statistical effects caused by the relatively small number of massive stars.  Some recent UV studies \\citep{Bouret2003,Fullerton2006} have called for a more fundamental mass-loss rate reduction, invoking clumping factors up to $\\sim$100, much higher than the currently favoured values of $\\sim 5$, and mass-loss rate reductions of order 10. However, other studies cast doubt on these conclusions based on theoretical studies of \"macro-clumping\" \\citep{Oskinova2007,Sundqvist2010} and emission in the extreme UV band \\citep{Waldron2010}. As we find good agreement between our population synthesis and the observations of Orion, this could either suggest that our mass-loss rates are realistic (and the very large clumping factors exaggerated), or alternatively that some unknown process is also missing in the stellar models. However, we note that the wind and supernova contributions to the interstellar $^{26}$Al have not been disentangled observationally, and models with weak winds to some degree compensate for the low $^{26}$Al wind yields by having larger core masses and therefore producing higher supernova yields \\citep[see discussion in][]{Limongi2006}, and the production of $^{26}$Al in high-clumping models have not yet been explored. On the other hand, we emphasize that our models are in \\textit{simultaneous} agreement with both the kinematic and the radioactive tracers, which would be hard to achieve with models involving very large clumping factors.  The Galaxy contains hundreds of regions of massive star-formation. It is important to extend our approach to other regions in order to overcome the issue of small-number statistics, and to further test our models. Unfortunately, many such regions are either significantly less well studied than Orion, due to larger distances and obscurance from the foreground, or they only contain modest numbers of high-mass stars. Recent studies of the relatively small, but nearby Scorpius-Centaurus region \\citep{Diehl2010} and the more distant but very massive Cygnus region \\citep{Martin2009} have been reported, and show overall agreement between the observations and models, both regarding the energetics of the regions and the $^{26}$Al signal. A further candidate target is the Carina region \\citep{Smith2006}, hosting a large population of very young and very massive stars. Due to the small age, the supernova contribution to the $^{26}$Al signal in this region is expected to be low. A comparison between Orion and Carina $^{26}$Al signal could therefore potentially constrain the relative wind and supernova contributions, similar to what has been done in the Cygnus region \\citep{Martin2009}.  Radioactive tracers are a valuable addition to the arsenal of probes of star formation in the Milky Way. $\\gamma$-ray observations have the potential to yield information that is complementary to observations at other wavelengths. They are emitted on a timescale of Myr after star formation, and with similar decay timescales, they trace the cumulative action of very young stars in the Milky Way. As the $\\gamma$-rays are unaffected by extinction, the observations of $^{26}$Al and $^{60}$Fe have the potential to give a complete view of the star formation in the Milky Way, unaffected by the obscuring effects of the molecular clouds. Using nucleosynthesis ejecta, we can expand the studies of past activity from stellar groups substantially. Instrument sensitivities of current $\\gamma$-ray telescopes limit such studies to the brightest, hence most-nearby regions. A next generation of instruments \\citep{Greiner2009} could reach hundreds of massive-star regions, thus significantly extending such studies. As discussed in \\citet{Voss-popsyn} the correlation between $^{26}$Al and $^{60}$Fe can potentially eliminate much of the uncertainty due to small-number statistics. The lifetimes of the radioactive tracers are long enough that they can be carried to significant distances from the massive stars that produced them. The COMPTEL observations of Orion \\citep{Diehl2002} provide a hint of such displacement. On the other hand the lifetimes are short enough that the radioactive elements are not uniformly mixed into the ISM. They are therefore valuable tracers of the mixing processes in the vicinity of star forming regions."
    },
    "1005/1005.1225_arXiv.txt": {
        "abstract": "Errors in numerical simulations of gravitating  systems can be magnified exponentially over short periods of time.  Numerical shadowing provides a way of demonstrating that the dynamics represented by numerical simulations are representative of true dynamics.  Using the Sitnikov Problem as an example, it is  demonstrated that unstable orbits of the 3-body problem can be shadowed for long periods of time.  In addition, it is shown that the stretching of phase space near escape and capture regions is a cause for the failure of the shadowing refinement procedure. ",
        "introduction": "The sensitivity which $N$-body integrations exhibit to small changes in initial conditions and to numerical errors has been an active area of research since Miller's landmark study.  \\cite{Miller} demonstrated the exponential divergence of near-by orbits for systems with $N\\leq 32$ and found that the separation of nearby orbits increases rapidly when close binary interactions occur.  He suggested that the divergence of near-by  orbits is too rapid to be solely accounted by binary interactions and suggests that there must be a collective effect to account for the results.  However, \\cite{Standish} showed that the divergence rate was reduced if the potential was replaced with a softened potential and concluded that the divergence is mainly due to close binary interactions.  The dramatic effects of numerical errors on $N$-body integrations was also demonstrated in an important paper by \\cite{Lecar}. After coordinating a study with 11 different integrations of the same 25-body problem for 2.5 crossing times, Lecar found that quantities such as half mass radius and the moment of inertia  can change by as much as 100 percent. In a study with $N=3$, \\cite{Hut} demonstrated that the amplification of initial errors can increase by as much as $10^{20}$.  In addition, they showed that the growth of errors during  close encounters can be amplified by as much as $10^4$, however some of the growth can be recovered  after the encounter is over.  The sensitivity to small changes in initial conditions and numerical errors is a property associated with chaotic systems. A measure of the sensitivity of numerical errors can be determined by the Lyapunov exponent $\\lambda$.  Early work suggested that the Lyapunov exponent is inversely  proportional   to the crossing time $t_{cr}$ \\citep{Kandrup,HeggieNbody,Goodman}. However, \\cite{Goodman} suggest a dependence on $N$ of the form $  \\lambda^{-1} = t_{cr}/\\log{N}$ or perhaps $\\lambda^{-1} = t_{cr}/\\log(\\log(N))$, implying that as $N$ increases  the rate of separation decreases and the Lyapunov exponent increases.  The $\\log(N)$ dependence was later numerically verified by \\cite{Merritt}.  Despite the difficulty calculating solutions to $N$-body integrations, computers still remain a useful tool to study self gravitating systems. If numerical errors in numerical solutions to the $N$-body problem cause such drastic changes in the actual positions and velocities of particles how can we trust the dynamics that these solutions represent? Shadowing is a way of proving that a true solution to a dynamical system follows close to a numerical solution.  If  true orbits can be found close to numerical orbits then the dynamics represented by the numerical solutions represents true dynamics.  This study will discuss the existence of shadow orbits for the gravitational 3-body problem. First,  definitions and concepts related to shadowing of dynamical systems will be introduced.  Next, a refinement procedure which makes corrections to  numerical orbits to reduce the errors incurred at each time step will be presented. The Sitnikov problem will then be presented and used as a simple model to discuss  escape and capture of orbits.  An approximate Poincar\\'e map is then presented  to model orbits of the Sitnikov problem and will be used in conjunction with the refinement procedure to discuss the validity of numerical solutions by way of shadowing. The failure of the refinement procedure to find shadow orbits will then be discussed and regions of phase-space where the procedure fails will be delineated. Finally, it will be demonstrated that the shadow times for this problem can be modeled as a Poisson process. ",
        "conclusions": "The above results confirm, for short lived orbits, previous investigations \\citep{Hayes} that showed numerical shadow durations, $M$, for gravitating systems follow a Poisson process with a exponential density function.  The result found in this study suggests for longer lived orbits, the density function is better approximated  by a function proportional to $1/M$.  This may be because the population of longer lived orbits tends to be dominated by stable orbits, however this has not been investigated.   In section \\ref{sec:5}, areas of phase-space where the refinement procedure is more likely to fail are characterized. These areas are near escape boundaries where  there is sufficient stretching of phase-space to cause the refinement procedure to fail to converge to a less noisy orbit.  Interestingly this seems to be due to the growth of the variational equations over one time step. This does not rule out the failure of the the refinement procedure by the accumulative effect of the growth of the variational equation associated with large Lyapunov exponents as discussed by \\cite{Zhu}.  In Figure \\ref{fig:ECCvsPERCENT} it is demonstrated that as the volume of phase-space representing capture orbits decreases, it becomes increasingly difficult to shadow capture orbits.  This is a result of the distribution of failures of the refinement procedure seen in Figure \\ref{fig:Edist}.   As the volume of  phase-space associated with capture decreases, capture orbits get pushed up against the boundary $\\partial \\mathcal{D}_0$ where the one-step growth of the variational equations causes the refinement procedure to fail.  Finally, it was found that the shadow distance for an orbit is proportional to the number of iterations of the map (Figures \\ref{fig:TvsEPSILON} and \\ref{fig:eps09}).  It was noticed that if in addition to $t_1$ and $E_1$, orbits were required to be shadow-able at the half steps $t_{1/2}$ and $E_{1/2}$, then initially shadow-able orbits continued to be shadow-able. When shadowing at the half step was required, the shadow distance typically increased by about a factor of two.   The Sitnikov problem discussed in this study provides a straight forward way of characterizing a domain of initial conditions as well as regions of stable and unstable motion. Work in progress considers slight changes to the Sitnikov problem in order  to study shadowing of unstable orbits.  For example, \\cite{Soulis1} consider slight perturbations to the mass and  position (away from the $z$-axis) of $m_3$ and delineate regions of stable and unstable motion. It would be expected that, like the results found in this study, shadowing with the refinement procedure breaks down near boundaries of escape for unstable orbits.   In fact, the break down of the refinement procedure near escape boundaries would be expected for  general 3-body configurations.  As solutions approach parabolic escape boundaries, an orbit can undergo increasingly  long ejections from the left-over binary system.  Small changes in the energy of an orbit in this region can cause significant changes in the time of return for the orbit. If the refinement procedure could make changes to the orbits so as to conserve the energy of the ejected body it might improve the success rate of the refinement procedure. Finally, the Sitnikov 4-body problem \\citep{Soulis2} provides a starting point for examining the relationship between the shadowing distance and the number of bodies.  Extra bodies can be added in circular orbits about the center of mass. \\cite{Hayes} demonstrates that as the number of moving bodies in a fixed potential increases the shadow durations decrease. It would be of interest to determine if a similar relationship holds for the Sitnikov $N$-body problem.   It should be stressed again that the failure of the  refinement procedure does not necessarily mean that a shadow does not exist for a given pseudo-orbit.   It may very well be that shadows do exist for orbits in regions where the refinement procedure fails.  We are encouraged that this may be the case.  Both the Sitnikov problem and the approximate Poincar\\'e map possess a hyperbolic invariant set, $\\Lambda$, near the escape boundaries (see \\cite{Moser} and \\cite{UrminskyThesis} respectively).  Despite the fact that $\\Lambda$ is near the boundary $\\partial \\mathcal{D}_0$, the shadowing theorems by \\cite{Anosov} and \\cite{Bowen} guarantee that any pseudo-orbit on $\\Lambda$ has an associated shadow-orbit.  This demonstrates that being in the vicinity on the escape boundary does not necessarily rule out the existence of shadow-orbits."
    },
    "1005/1005.3220_arXiv.txt": {
        "abstract": "{ We explore the kinematical effect of having extra dimensions on the gravitational wave emission from cosmic strings. Additional dimensions both round off cusps, and reduce the probability of their formation. We recompute the gravitational wave burst, taking into account these two factors, and find a potentially significant damping on the gravitational waves of the strings. } ",
        "introduction": "\\label{intro}  The notion that nature might have extra dimensions, and that these might leave traces in observation or experiment, has led to many interesting ideas and advances over the last decade. Traditionally, extra dimensions would be hidden via a Kaluza-Klein mechanism, accessible only at ultra high energies. More recently however, the braneworld paradigm, \\cite{EBW,ADD,RS}, has allowed for extra dimensions to be much larger, hidden instead by a confinement mechanism, \\cite{EBW}, which localizes standard model physics on the brane but allows gravity to probe the bulk. These Large Extra Dimension (LED) scenarios, \\cite{ADD,RS}, have the added attraction of providing a natural hierarchy between gauge and gravity interactions coming from geometric multiplying factors in the derived 4-dimensional Planck mass. These ideas have been incorporated into string theory models, allowing LED's via a process of flux stabilization, \\cite{KKLT}.  One of the advantages of the braneworld scenarios was that they gave concrete predictions for cosmology. In particular, the Randall-Sundrum (RS), \\cite{RS} set-up, with a brane living in anti de Sitter (adS) spacetime with one extra dimension, has a simple cosmology, the scale factor being determined by motion through a black hole spacetime, \\cite{BCOS}. This led to the mirage picture, \\cite{mirage}, in which brane cosmology is determined by bulk motion through a warped space, however, the mirage picture does not explicitly include the gravitational back reaction of the braneworld, which for codimension two and higher can be problematic, \\cite{HCoD}.  Nonetheless, the idea of the braneworld -- an object localized in extra dimensions -- has proven to be extremely useful when constructing models of inflation in string theory. The localization in extra dimensions provides a natural scalar field (the position of the brane) which can then evolve according to some effective action depending on the specifics of the internal manifold and the brane itself. Early models of brane inflation, \\cite{braneinf}, used the Newtonian potential of the higher dimensional solution to provide the inflaton potential. This however, was not consistent with the size of the extra dimensions, nor was any mechanism offered for their stabilization. However, with the completion of the picture via flux stabilization, \\cite{KKLT,KKLMMT}, a great many models of brane inflation have been explored (see \\cite{BI} for a review).  A key side effect of brane inflation is the formation of cosmic strings, \\cite{BCS,JST}. These strings form as a by product of the annihilation of the inflationary branes (for reviews see \\cite{CSSrev}), and while having their origins in superstring theory, can have a wide range of parameters, and interesting physical properties.  Cosmic strings, \\cite{CS}, are an example of a topological defect, a glitch in the vacuum structure of a field theory which can arise when the vacuum manifold is topologically nontrivial. They are ubiquitous in all types of physics, from condensed matter systems to quantum field theory. They were first explored in the cosmological setting by Zeldovich and Kibble, \\cite{ZelKib}, who realised that they could arise in the early universe as possible side-products of symmetry breaking phase transitions. Indeed, it was the cosmological catastrophe of monopole production at the GUT scale that in part led to the development of the original inflationary scenario. Cosmic strings however are cosmologically benign, and exhibit a scaling network behaviour, \\cite{NET}. Accounting for their gravitational effects made them a possible candidate for the perturbation spectrum until CMB experiments ruled them out, \\cite{CMBrout}. Nonetheless, the possibility remains that a network of light strings could be present in our universe, \\cite{CSfit}.  From the cosmological point of view, the internal structure of the cosmic string is irrelevant, and what we need is the long range behaviour of the string. Gravitationally, a straight string produces a conical deficit in spacetime \\cite{cone}, which, while it does lead to interesting lensing effects, cannot be regarded as a detection tool as it requires the serendipity of an appropriately aligned source behind the string (although see \\cite{KK} for a discussion of detection via weak lensing). Instead, a more promising approach is to take the dynamical network of strings, and to use linearized gravity to compute the radiation emitted from the loops and crinkly long strings \\cite{gwave,GWL,Eloss}. Network simulations approximate the string by the Nambu action: \\be \\label{NGaction} S = -\\mu \\int d^2 \\sigma \\sqrt{\\gamma} \\ee which can be rigorously derived from an underlying field theory model \\cite{NG}. Together with rules for intercommutation \\cite{IC}, or how strings behave when they cross each other, this gives the basic physics of how a network evolves.  The primary drivers of network evolution are therefore whether strings intersect, and whether the motion has any extreme events (i.e.\\ are there points at which the Nambu approximation is likely to break down). Early work focussed on taking toy families of loops, \\cite{KT,loops}, to explore the likelihood of self intersection, and also events where the string bends on the same scale as its width. The kinematics of strings obeying the Nambu action are very well understood: the picture is that loops often self intersect, and generically have cusps -- points at which the string reaches the speed of light and its extrinsic curvature diverges, leading to a breakdown of the Nambu action. Finite width corrections to the Nambu action can be computed, \\cite{FWC,EffAction}, but it was believed that these would not be significant for the network evolution. In appendix \\ref{AppA} we demonstrate this explicitly by computing the finite width effect for a cusp. More recently, further analysis of the effect of small scale structure on the string, \\cite{SSS}, has indicated that there may actually be measurable consequences. In either case however, the primary importance of a cusp is that it acts as a strong source of gravitational radiation.  In a seminal paper, Damour and Vilenkin (DV), \\cite{DV}, re-examined gravitational radiation from cosmic strings, assessing for what range of mass per unit length the string could potentially be visible to the next range of gravitational wave detectors. The main effect they were considering was the burst of radiation from extreme kinematic events in the loop motion, known as cusps and kinks. They computed the amplitude of a cusp and kink gravitational wave burst (GWB) as a function of the mass per unit length of the string. In a later paper, \\cite{DV2}, they allowed for networks formed by strings with lower intercommutation probabilities, which enhance the density and thus the GWB amplitude. Siemens et al.\\ \\cite{Siem,SiemStoch}, performed a more careful analysis of the cosmological expansion history, instead computing rates of events at amplitudes fixed by the detector. Since then, many other gravitational effects of cosmic superstrings have been explored, including strings with junctions, \\cite{GWCSS}, and broken strings, \\cite{beads}.  In a recent note \\cite{CCGGZ}, we revisited the calculations of DV et al.\\ arguing that the kinematic effect of the extra dimensions significantly reduced the power of the cusp waveform. Essentially, extra dimensions act to `slow down' the string, as first pointed out by Avgoustidis and Shellard, \\cite{AS}, and round off the sharp cusp. In this paper we give a full computation of this effect, presenting test loop families to demonstrate parameter space measures, and detailed numerical calculations of GWB event rates and amplitudes.  The extra dimensions give two main modifications to the DV result. The first is the cusp rounding effect, and we show how this gives a high frequency cut-off to the gravitational waveform. The second is a probabilistic factor: in $3+1$ dimensions, cusps always form on a smooth loop trajectory, however, in higher dimensions this is no longer the case. DV introduced a parameter in their GWB calculation, ${\\cal C}$, which measured the probability of cusp formation in a single loop period; the reason for introducing this in $3+1$ dimensions is that once strings intersect, they can have kinks, which lower the cusp probability. In higher dimensions however, the probability of cusp formation is strictly zero, in that the set of solutions which have cusps has measure zero. We therefore define a {\\it near cusp event}, which represents the rounded cusp, and compute the probability of this event to input in the amplitude calculation.  Finally, one has to factor in the reduction in intercommutation probability due to the extra dimensions, \\cite{ICSC,AS2}, which acts to increase the density of the network. Summing all these effects produces a marked effect on the GWB amplitude, and our conclusion is that detection of GWB's from cosmic superstrings by current or next generation gravitational wave detectors will be harder than suspected, with the bonus that positive detection may tell us something about the number of extra dimensions. Clearly, our result will also relax bounds on $G\\mu$ for cosmic superstrings derived using the DV results \\cite{obsv}.  We start by reviewing the standard Kibble-Turok method of analysing string trajectories, \\cite{KT}, noting the new features that appear with additional dimensions. We then review the DV calculation of the GWB. Next we calculate the various effects coming from extra dimensions, introducing a parameter $\\Delta$ which measures deviation from an exact cusp. We integrate over $\\Delta$ to obtain the sum of all cusp or near cusp events, and finally discuss implications and caveats of the calculation. To be specific, we focus on the frequency band of the advanced LIGO detector, however we will comment on the frequency dependence of our results. ",
        "conclusions": "\\label{disc}  The clear message of our results is that the impact of motion in extra dimensions can be significant. That extra dimensions should have an effect is not unreasonable, since they can be viewed as additional degrees of freedom living on the string\\footnote{We thank Jose Blanco-Pillado for discussions on this point.}, for example, superconducting cosmic strings, \\cite{Wit}, can be represented as a dimensional reduction of standard five dimensional KK theory \\cite{SCS}. For these cosmic strings, the currents round off cusps in much the same way as we have described here, \\cite{scusp}, and also alter the balance between the energy and tension of the string, \\cite{SCgrav}, which has a clear gravitational impact. This naturally raises the concern that we may not be able to distinguish between extra fields living on a cosmic string, and extra dimensions in which the string is moving. This would certainly be the case if one was observing a string and a single GWB. However, our calculation was for the expected signal from a cosmological network of strings, which depends not only on the GWB waveform, but also on the properties of the network. Superconducting string networks have not been as well explored as those of standard cosmic strings, \\cite{scnet}, with the main focus being on the different physics induced by the long range electromagnetic interactions. Nevertheless, as superconducting strings have similar intercommutation properties to standard cosmic strings \\cite{sci}, it is likely that the network is more similar to the usual cosmic string network than that of the cosmic superstrings. Thus, while the individual GWB waveforms will be similar, the expected rates and signals we have calculated for the cosmological networks are indeed specific to extra dimensions.  It is perhaps a little surprising that the effect of extra dimensions can be so dramatic. We therefore now examine our assumptions carefully, raising below a series of critiques together with a discussion of their validity and resolution.  The basic reason for the suppression of the signal is the distribution over the near cusp parameter $\\Delta$. Our simulations with test loop trajectories were performed in flat space with a toroidal (flat) compactification. Clearly the cosmological evolution will influence the distribution of momentum across the various modes (although paradoxically it would seem to damp higher momenta in the noncompact dimensions). Since these strings form from the collision of a brane and anti-brane, it seems likely that they have significant initial momentum in the extra dimensions, thus we see no reason to curtail our solution space in this way. The main objection to having total freedom of the internal modes is that by wrapping back and forth across the extra dimension(s) the string has more opportunity to self intersect, thus curtailing the additional freedom in that direction.  We modelled this effect by exploring the self-intersection of a $1-3/1$ family of loops. We chose this combination of harmonics, as the 4D family will have a 3D limit which were the first simple loop trajectories explored by Kibble and Turok \\cite{KT}. In 3D, the loop family self intersects approximately $30\\%$ of the time (note, the original plot of Turok \\cite{KT} is inaccurate, see \\cite{CDH} for the correct version). When an additional dimension is introduced, the measure of the solution set allowing for self intersections again becomes zero: the argument is once more parametric. A string will self-intersect if \\be \\Ab(\\tau-\\sigma) + \\Bb(\\tau+\\sigma) = \\Ab(\\tau-\\sigma') + \\Bb(\\tau+\\sigma') \\label{selfint} \\ee for some $\\tau, \\sigma, \\sigma'$. Thus, there are three dynamical variables and $3+n$ constraints. In 3D, it is therefore possible to satisfy these constraints simultaneously, although a more careful check of the parameters shows that not all loop solutions can satisfy these constraints. Nonetheless, it shows how the subspace of self intersecting loops can be of nonzero measure in parameter space. With extra dimensions however, satisfying (\\ref{selfint}) necessarily requires a constraint on parameter space, hence the subspace with self intersections will be of lower dimension than parameter space. Even compactifying the extra dimensions does not change this argument, unless we take into account the finite width of the string. Essentially, if we take the string to have zero width, then it can easily miss itself even when winding back and forth across the extra dimensions many times. However, with finite width, the self intersection probability now becomes nonzero, and of order ${\\cal O}(w/R)^n$ (where $w$ is the string width, and $R$ the size of the extra dimension). This therefore suggests that we introduce this ratio in a finite width correction to the GWB measure.  \\vskip 3mm  As we mentioned during our initial discussion, warping of the extra dimensions could also provide some significant dynamical effect. The results of Avgoustidis \\cite{AA} indicate that warping is not as dramatic a trapping force as intially suspected, however, any confinement of strings could be significant, and a detailed modelling of this effect is necessary. For now, we model a restriction in the extra dimensional motion by a restriction in $\\Delta$. Instead of allowing $\\Delta$ to range over the full unit interval, we restrict $\\Delta\\in[0,\\Delta_0]$. Thus we must renormalize ${\\cal C}$: \\be {\\cal N}(\\Delta_0) = \\int_0^{\\Delta_0} {\\cal C}(\\Delta) = 1 \\qquad \\Rightarrow \\qquad {\\cal C}(\\Delta) = \\frac{n}{\\Delta_0^n} \\Delta^{n-1} \\;. \\label{delta0C} \\ee Thus the relevant $\\Delta$ integral (\\ref{deltaint}) now becomes \\be \\label{deltaint0} \\int_0^{{\\rm min}\\{\\Delta_0,\\theta_m\\}} {\\cal C}(\\Delta) \\left ( \\theta_m(z) - \\Delta \\right )^2 = \\theta_m(z)^2 {\\cal F}_n \\left [ \\frac{\\theta_m}{\\Delta_0} \\right ] \\ee where \\be {\\cal F}_n [x] = \\frac{2 x^n}{(n+1)(n+2)} H[1 - x] + \\left (1 - \\frac{2n}{(n+1)x} + \\frac{n}{(n+2) x^2} \\right ) H[x-1] \\;. \\ee  A reasonable value for $\\Delta_0$ might be to use the one parameter ratio that does impact on the loop families and motion: the ratio of string width to the size of the extra dimension. The Nambu action is only a good approximation when the width of the string is small compared to scales of physical interest. This is rarely a problem in cosmology, as the string width is set by the inflationary scale, and the size of the universe rapidly becomes large. As far as the extra dimensions are concerned however, these are stabilized at a couple of orders of magnitude above the string scale, hence while the Nambu action is a good approximation, we might expect some corrections to show up due to parameter restriction from self intersection or excessive winding as already discussed. We therefore expect this parameter to be related to the probability of intercommutation, which can be viewed as arising because of the strings ``missing'' each other in the internal dimensions.  To test this alternate expression, we took values of $\\Delta_0 = 0.1, 0.01, 0.001$, and $10^{-4}$. These values were motivated by a limiting sensible ratio $w/R$, and the value of $P$. From (\\ref{deltaint0}), we see that the effect of $\\Delta_0$ is to cut off the integral as $\\theta_m$ grows. For $\\theta_m < \\Delta_0$, the dependence of the rate on $\\theta_m(z)$ remains that of the previous section, however, as $\\theta_m$ grows, the functional dependence shifts towards the $\\theta_m(z)^2$ form of the 3D result. From the expression for $\\theta_m(z)$, (\\ref{thetaofz}), we see that this is proportional to $(G\\mu)^{-1/3}$, hence the rates converge to the 3D value sooner for smaller $G\\mu$. \\FIGURE{\\label{fig:RateDel} \\includegraphics[width=14cm]{FRate21Dl.eps} \\caption{ A plot of the dependence of the rate on $\\Delta_0$ with the number of extra dimensions fixed at three. From top to bottom, the solid black line indicates the 3D result, $\\Delta_0 = 10^{-4}$ in dashed green, $\\Delta_0 = 10^{-3}$ in dot-dash red, $\\Delta_0 = 10^{-2}$ in dashed purple, and $\\Delta_0 = 10^{-1}$ in dotted blue respectively. The horizontal black line indicates a rate of one event per year. } }  Figure \\ref{fig:RateDel} shows the effect of the $\\Delta_0$ parameter on the event rate at an amplitude cutoff of $10^{-21} s^{-1/3}$ for $\\Delta_0$ ranging from $10^{-1} - 10^{-4}$ as indicated. Note that once we use this more complicated expression (\\ref{deltaint0}), the use of the interpolating function approximation becomes too unwieldy, and the rates had to be calculated by direct integration. Figure \\ref{fig:Raten} shows the effect of the rate on the number of extra dimensions, fixing $\\Delta_0 = 10^{-3}$ and allowing $n$ to vary as indicated. Here we see that for all $n$ the plots converge on the 3D result at $\\alpha \\sim 10^{-11}$ but for $\\alpha \\sim 10^{-8}$ for example, the rate drops by roughly an order of magnitude per dimension. A positive detection therefore could potentially tell us the number of extra dimensions! \\FIGURE{\\label{fig:Raten} \\includegraphics[width=14cm]{FRate21nn.eps} \\caption{ A recomputation of the rate plot with the new nearcusp measure (\\ref{deltaint0}) fixing $\\Delta_0 = 10^{-3}$, and allowing $n$ to vary. From top to bottom, $n=0, 1, 3$ and $6$ respectively. The horizontal black line indicates a rate of one event per year. } }  To sum up: We have included the effect of motion in extra dimensions in the computation of the gravitational wave burst signal from cusp events on cosmic string loops. We find a significant effect, even after taking into account finite width effects and the size of the extra dimension. Clearly further work is required to get better control of the approximations being used, in particular to take into account more complex compactification geometries, however it does seem that motion in internal dimensions is important. Finally, if the string tension lies in the serendipitous range $10^{-8}-10^{-10}$ then the possibility arises that a positive detection by gravitational radiation would not only confirm the general brane inflation scenario, but could provide a means of determining the number of (effective) extra dimensions."
    },
    "1005/1005.0709_arXiv.txt": {
        "abstract": "\\src\\ is a supergiant high mass X-ray binary discovered by INTEGRAL in 2004. The source is detected at a significance level of $\\sim21$ standard deviations in the \\sw-BAT survey data collected during the first 54 months of the \\sw\\ mission. The timing analysis reveals an orbital period of $\\sim$6.78 days and the  presence of a full eclipse of the compact object. The dynamical range (variability up to a factor $\\sim$20) observed during the BAT monitoring suggests that \\src\\ is a wind-fed system. The derived semi-major axis of the binary system is $\\sim55 R_{\\sun}$ with an orbit eccentricity lower than 0.15. ",
        "introduction": "}  The IBIS/ISGRI telescope \\citep{Ubertini03,Lebrun03} on board the INTEGRAL satellite \\citep{Winkler03} has detected a large number of new supergiant high mass X-ray binaries (sgHMXB) characterized by a high absorption column (N$_{\\rm H}> 10^{23}$ cm$^2$) and/or by fast bright transient events. This result was achieved thanks to a continuous monitoring of the Galactic plane with deep exposure of the Galactic centre and thanks to a combination of ISGRI large field of view, good sensitivity, and wide energy range. The fourth IBIS/ISGRI survey catalogue \\citep{Bird10} reports the discovery of 32 new objects that have been classified as HMXB and  84 unidentified sources with a Galactic latitude $|b| < 5^{\\circ}$.   The Burst Alert Telescope (BAT, \\citealp{Barthelmy2005:BAT}) on board \\sw\\ \\citep{Gehrels2004mn} has been performing a continuous monitoring of the sky in the hard X-ray energy range (15--150 keV) since November 2004. The telescope,  thanks to its large field of view (1.4 steradian half coded) and its pointing strategy, covers a fraction between 50\\% and 80\\% of the sky every day. This has allowed the detection of many of the new INTEGRAL HMXBs \\citep[e.g.][]{Cusumano2010} and the collection of their long term light curves and of their spectral energy distributions. The long and continuous monitoring of these sources allows to investigate the intrinsic emission variability, to search for long periodicities (orbital periods) and to discover the presence of eclipse events. The role of \\sw-BAT is therefore fundamental to unveil the nature and the geometry of these binary systems.  In this Letter we analyze the soft and hard X-ray data collected by \\sw\\ on IGR~J16493--4348. This source was discovered by INTEGRAL in 2004 \\citep{Bird2004} and it was initially associated with the radio pulsar PSR~J1649--4349 because of a spatial coincidence. A later INTEGRAL observation with a deep exposure allowed to reduce the positional uncertainty and to reject the pulsar association \\citep{Grebenev2005}. A follow-up observation with {\\em Chandra} found a soft X-ray counterpart at RA(J2000) = 16$^{h}$ 49$^{m}$ 26.92$^{s}$; Dec(J2000) = -43$^{\\circ}$ 49' 8.96'' \\citep{Kuiper2005} allowing the optical association with 2MASS~J1642695--4349090, a B0.5 Ib supergiant \\citep{Nespoli2008}. The {\\em Chandra} observation also revealed strong evidence for variability. The IGR~J16493--4348 energy distribution extracted from two  RXTE observations \\citep{markwardt05} was modeled with a highly absorbed power law (N$_{\\rm H}\\sim10^{23}$ cm$^{-2}$)  with a photon index of $\\sim$1.4 and a 20--40 keV flux of 2.1$ \\times 10^{-11}$ \\ferg. A spectral analysis using non simultaneous data from \\sw-XRT and INTEGRAL \\citep{hill07} showed the presence of a cutoff at $\\sim15$ keV. The information derived from the X-ray observations and the identification of the spectral type of the optical counterpart allowed the classification of \\src\\ as a supergiant High Mass X-ray Binary (sgHMXB).  This Letter is organized as follows. Section 2 describes the data reduction. Section 3 reports on the timing analysis and in Sect. 4 we discuss our results. ",
        "conclusions": "}  We have presented the results obtained from the analysis of the data collected by \\sw-BAT during the first 54 months of the Swift mission on the supergiant HMXB IGR~J16493--4348. The source is detected at a significance level of $\\sim$21 standard deviations.  The long monitoring coupled with a good observation duty cycle ($\\sim13\\%$ per day) allowed us to unveil a periodicity of $\\sim6.782\\pm0.002$ d  that we interpreted as the orbital period of the binary system.  The knowledge of the orbital period allows us to derive the the semi-major axis of the binary system through the Kepler's third law: $a^3=G$ P$_0^2$$(M_{\\star}+M_{\\rm X})/4\\pi^2$ where $M_{\\star}$ and $M_{\\rm X}$ are the masses of the supergiant and compact object, respectively. Adopting $M_{\\star}=47\\,M_{\\sun}$ \\citep[][for an B0.5 Ib star, of radius $R_{\\star}=32.2\\,R_{\\sun}$]{Searle2008} and a typical neutron star mass of $M_{\\rm X}=1.4\\,M_{\\sun}$ we obtain $a\\sim55\\,R_{\\sun} \\sim 2\\,R_{\\star}$. This orbital separation is  common among classical sgHMXB \\citep{Walter2007}. The folded light curve shows the presence of a full eclipse with an estimated time of mid-eclispe at T$_{\\rm eclipse}=(54175.92 \\pm 0.10) \\pm  n \\times$ P$_0$ MJD and it appears flat outside the eclipse. The duration of the eclipse ($\\sim 12\\%$ of the orbital period) is roughly consistent with the occultation of the compact source by the companion star in a highly inclined orbit with low eccentricity.  The 15--50 keV light curve of IGR~J16493--4348 is variable up to a factor $\\sim$20 on timescales of $\\sim$300 s while its average over an orbital period shows a variability up to a factor $\\sim$4. The inferred orbital separation suggests a wind-fed system. This hypothesis and the radius of the supergiant allows to limit the eccentricity of the orbit.  Figure~\\ref{lcc}  shows how the Lagrangian point  between the two stars changes as a function of the orbital phase for different orbit eccentricities \\citep{paczynski71}; the maximum eccentricity compatible with a wind-fed system (Lagrangian point lower than the supergiant radius) is 0.15.     \\begin{figure}% \\begin{center} \\centerline{\\includegraphics[width=6cm,angle=-90]{figure4.ps}} \\caption[]{The Lagrangian point of the binary system as a function of the orbital phase for different orbital eccentricity. The straight dashed line represents the radius of the companion star. } \\label{lcc} \\end{center} \\end{figure}"
    },
    "1005/1005.2154.txt": {
        "abstract": "%We aim at showing that elliptical galaxies formed through dissipative %collapse %exhibit a range of radial gradients in the stellar metallicity %comparable to %the observed one. { We investigate whether it is possible to explain the wide range of observed gradients in early type galaxies in the framework of monolithic models}. %The main scope is to understand whether a correlation %of the gradient slope with the galactic mass exists. To do so, we extend the set of hydrodynamical simulations by Pipino et al. (2008a) by including low-mass ellipticals and spiral (true) bulges. These models satisfy the mass-metallicity and the mass-[$\\alpha$/Fe] relations. The typical metallicity gradients predicted by our models have a slope of -0.3 dex per decade variation in radius, consistent with the mean values of several observational samples. However, we also find a few quite massive galaxies in which this slope is -0.5 dex per decade, in agreement with some recent data. In particular, we find a mild dependence from the mass tracers when we transform the stellar abundance gradients into radial variations of the $Mg_2$ line-strength index, but not in the \u00e5\u00ca$Mg_b$. We conclude that, rather than a mass- slope relation, is more appropriate to speak of an increase in the scatter of the gradient slope {with the galactic mass}. %at a fixed mass. We can explain such a behaviour with different efficiencies of star formation in the framework of the revised monolithic formation scenario, hence the scatter in the observed gradients should not be used as an evidence of the need of mergers. %rather than by splitting galaxies into several %categories according %to their formation mechanism; namely truly monolithic ones versus %merger outcomes, {as recently suggested in the %literature Indeed, model galaxies that exhibit the steepest gradient slopes are preferentially those with the highest star formation efficiency at that given mass. ",
        "introduction": "Negative metallicity radial gradients in the stellar populations are a common feature in spheroids (e.g. Carollo et al. 1993, Davies et al. 1993) and must be predicted by every theory for the formation of elliptical galaxies. A possible fingerprint of a given galaxy formation scenario might be the (lack of) correlation between gradient properties (e.g. the slope) and either global galactic properties (namely mass, stellar velocity dispersion $\\sigma$, total magnitude) or central ones (e.g. central metallicity or [$< \\alpha$/Fe$>$]). From the theoretical point of view, in fact, steep metallicity gradients are expected from classical dissipative collapse models (e.g. Larson 1974, Chiosi \\& Carraro, 2002) and their (\\emph{revised}) up-to-date versions  which start from semi- cosmological initial conditions (e.g. Kawata 2001, Kobayashi, 2004). The abundance gradient arises because the stars form everywhere in a collapsing cloud and then remain in orbit with a little inward motion,\\footnote{Stars will spend most of their time near the apocentre of their orbit.} whereas the gas sinks further in because of dissipation. This sinking gas contains the new metals ejected by evolving stars so that an abundance gradient develops in the gas. As stars continue to form their composition reflect the gaseous abundance gradient. { The original dissipative models} predict a steepening of the gradient as the galactic mass increases, mainly because the \\emph{central} metallicity is quickly increasing with mass \\footnote{The fit of the mass-metallicity relation, namely the increase of the mean metal content in the stars as a function of galactic mass (O'Connel, 1976), was the main success of these {} original models.}, whereas the \\emph{global} one has a milder variation (Carlberg 1984). { At the same time they} predict metallicity gradient as steep as -0.5 dex per decade variation in radius. { On the other hand, the few attempts to study the gradients in the merger-based models hint for very shallow (if any) gradient (Bekki \\& Shyoia, 1999), less steep than the mean observational values and than the predictions from monolithic collapse models. Moreover, it seems that dry mergers flatten pre-existing gradients (Di Matteo et al., 2009). Indeed, when the two scenarios (monolithic collapse and mergers) are considered as two possible channels working at the same time, the scatter in the predicted gradients for such a population of galaxies seem to be in agreement with observations (Kobayashi, 2004).}  %At variance with %earlier monolithic collapse models, the chemical evolution model by %PM04 takes into account a gas infall %and a star formation timescale both decreasing as a function of the %galactic mass, in order to satisfy the [$<\\alpha$/Fe$>$]-mass\\footnote{We will %always refer to average %abundances and abundance ratios by including the symbols $<\\, ,\\, >$ %in the standard %notation , being these values closer to the SSP-equivalent value %inferred from observations for composite stellar populations (see %Pipino et al., %2006 for details, hereafter PMC06). The standard notation, i.e. [$ %\\alpha$/Fe], will be used %for the gas abundance ratios only.} %relation (Worthey et al. 1992, see Matteucci 1994 for a detailed %discussion). %PM04 predicted a typical gradient slope of -0.25 dex per decade %variation in radius %and a flat trend of the metallicity gradient slope with galactic mass, %in agreement with the observations available at that time (e.g. %Kobayashi \\& Arimoto, 1999).  %The shallower gradient predicted by this new monolithic %models with respect to earlier dissipational collapse models as well %as the lack of correlation between gradient slope and galactic mass %stem from two reasons: i) no gas flows are allowed %between regions at different galactic %radii (i.e. no strong dissipation {contrary to what %happens} in the Carlberg (1984) models); ii) {we  must account %for a constraint that was not available when the original collapse models %were put forward, namely it is necessary to produce stars with high [$<\\alpha %$/Fe$>$] ratios inhabiting the galactic core. This turns into a need %for short duration of the star formation. Hence the %metal enrichment process in the centre is faster, leading to smaller %age gradients} even in the most massive galaxies. %A clear prediction of PM04 model was a positive gradient in the [$< %\\alpha/Fe>$], arising from the outside-in formation of ellipticals, %where the star formation stops first in the outermost regions {(see also Pipino et al. 2006)}.  More recently, observations showed that successful models for elliptical galaxies should also reproduce the [$<\\alpha$/Fe$>$]-mass relation (Worthey et al. 1992, Thomas et al. 2007) as well as the observed %This result seems to be questioned by a handful of more recent observational %works {which } gradients in the [$<\\alpha/Fe>$] ratios (Mehlert et al. 2003, Annibali et al. 2007, Sanchez-Blazquez et al. 2007, Rawle et al., 2008). {Indeed, these observations} show that the slope in the [$<\\alpha/Fe>$] gradient has a typical value close to zero and does not correlate with mass. %{This suggests that the outside-in formation scenario for ellipticals galaxies is %not a sufficient condition for the formation of the abundance ratio %gradients (see Pipino et al. 2008a, hereafter Paper I)}.  These observations have been interpreted by Pipino et al. (2008a; Paper I, hereafter) 1D hydrodynamical code as the fact that the suggested outside-in mechanism for the formation of the ellipticals is not the only process responsible for the formation of gradients in the abundance ratios. Other processes should be considered such as the interplay between the star formation (SF) timescale and gas flows. While such an interplay flattens the [$<\\alpha/Fe>$] gradient to the value required by observations, it still enables galaxies to harbor gradients in [$<$Fe/H$>$] and [$<$Z/H$>$] in agreement with the most recent observations (see Section 2). %Here we note that the shallower gradient predicted by this new monolithic %models with respect to earlier dissipational collapse models as well %as the lack of correlation between gradient slope and galactic mass %stem from the fact that it is necessary to produce stars with high [$<\\alpha %$/Fe$>$] ratios inhabiting the galactic core. %Such a constraint, that was not available when the original collapse models, %turns into a need %for short duration of the star formation. Hence the %metal enrichment process in the centre is faster, leading to smaller %age gradients} even in the most massive galaxies. Pipino et al. (2008b; Paper II) calibrated such a model by means of the resolved stellar populations in the Milky Way bulge. As a matter of fact, spiral \\emph{true}\\footnote{In the rest of the paper we will consider only the class of \\emph{true} bulges, (Kormendy \\& Kennicutt, 2004).} bulges remarkably follow many fundamental constraints for ellipticals such as the mass-metallicity and the mass-[$<\\alpha$/Fe$>$] relations (see below), the only difference being that they might be rejuvenated systems (Thomas \\& Davies, 2006).  The aim of this paper is to explore a wider range of cases by extending the analysis of Paper I to lower masses, including bulges, and compare them to the latest observational results. { In this way we can study the correlation between gradient slopes and galactic mass (if any) in order to understand whether the monolithic galaxy formation scenario is in agreement with the recent observational evidences}.  In Sec.~\\ref{obs} we give a brief overview of the observations regarding metallicity gradients in ellipticals. The main characteristic of the model are briefly described in Section~ \\ref{model}. We characterize the global properties of our models in Section 4, present our results in Section~\\ref{results}, discuss them in Section 6 and draw our conclusions in Section~\\ref{concl}. ",
        "conclusions": "\\label{concl}  In this paper we study the formation and evolution of ellipticals and bulges by means of a hydrodynamical model (c.f. Paper I and II, respectively) in order to understand the origin of the observed scatter in the abundance gradients of early type galaxies. Here we summarise our main results.  \\begin{itemize} \\item We find $\\Delta_{Z/H}$ in the range -0.5 -- -0.2 dex per decade in radius with a mean value of -0.3 dex per decade in radius, in agreement with the observations (e.g. Kobayashi \\& Arimoto, 1999). \u00e5\u00ca  \\item In agreement with Ogando et al. (2005) and Jablonka et al. (2007), we find that that the scatter in the gradient slopes increases as a function of mass. We reproduce such a scatter in the observations by means of variation in the initial conditions in galaxy models.  \\item Model galaxies which behave as the earlier \\emph{monolithic collapse} models by Larson (1974) and Carlberg (1984) define a \\emph{maximum steepness boundary} in the metallicity (and index) gradient slope-mass plane. \\emph{These galaxies are preferentially those with the highest star formation efficiency at that given mass.}  \\item No galaxies with gradients steeper (i.e. more negative) than the value given by the our predicted \\emph{theoretical boundary} are observed (Ogando et al., 2005, and Spolaor et al., 2009, for ellipticals and Jablonka et al., 2007, for bulges).  \\item No correlation between $\\Delta_{O/Fe}$ and other galactic properties are found, in agreement observations for ellipticals (Mehlert et al., 2003, Annibali et al., 2007) and bulges (Jablonka et al., 2007).  \\item The abundance gradients in the abundances, once transformed into line-strength indices, lead to $d Mg_2/log (R_{core,*}/R_{eff,*})\\sim -0.06$ mag and $d Log Mg_b/log (R_{core,*}/R_{eff,*})\\sim -0.1$ \u00e5\u00ca per decade in radius, again in agreement with the typical mean values measured for ellipticals and bulges.  \\item We note that the behaviour of the gradient slope as a function of the galactic mass strongly depends on the particular line-strength index used. In fact, the predicted $Mg_2$ index gradient seems to correlate with mass, whereas the $Mg_b$ index gradient does not.  \\item In Paper I we demonstrated that the differential occurrence of galactic winds (outside-in formation) alone can explain the existence of the metallicity gradients discussed in this paper. Here we add that this mechanism predicts a tight correlation between line-strength index and escape velocity gradients which has been confirmed by recent data (see Scott et al., 2009). \\end{itemize}   Larger, homogeneous and statistically meaningful observational sample of gradients in elliptical galaxies out to one effective radius can confirm such a prediction and validate the model.  \\clearpage  \\begin{figure} %\\epsscale{.80} %\\includegraphics[width=8cm,height=8cm]{fig2.eps} \\includegraphics[width=\\textwidth,height=15cm]{fig2_new_new.eps} \\caption{Comparison between model predictions and data in the observer's plane. \\emph{Upper panels}: Predicted $Mg_2$ gradient slope as a function of the stellar mass and mass tracers. The hatched area represents the portion of the plane $\\Delta Mg_2$-mass (-$\\sigma$) \\emph{allowed} by the data from Ogando et al. (2005). %Crosses: %data from Mehlert et al. (2003, red), Annibali et al. (2007, blue) %and Jablonka et al. (2007, yellow). Symbols and data as in Fig.~\\ref{fig1}. \\emph{Lower panels}: As in the upper panel, but for the $Mg_b$ index.} \\label{fig2} \\end{figure}  \\clearpage  %\\acknowledgment"
    },
    "1005/1005.5406_arXiv.txt": {
        "abstract": "Advances in stellar interior modeling are being driven by new data from large-scale surveys and high-precision photometric and spectroscopic observations.  Here we focus on single stars in normal evolutionary phases; we will not discuss the many advances in modeling star formation, interacting binaries, supernovae, or neutron stars.  We review briefly:  1) updates to input physics of stellar models; 2) progress in two and three-dimensional evolution and hydrodynamic models; 3) insights from oscillation data used to infer stellar interior structure and validate model predictions (asteroseismology).  We close by highlighting a few outstanding problems, e.g., the driving mechanisms for hybrid $\\gamma$ Dor/$\\delta$ Sct star pulsations, the cause of giant eruptions seen in luminous blue variables such as $\\eta$ Car and P Cyg, and the solar abundance problem. ",
        "introduction": "Advances in stellar modeling are being driven and validated by stellar oscillation data.  Over the last twenty years, multiple oscillation modes of many star types in all evolutionary phases have been discovered and observed from ground-based networks and satellites.  These modes include gravity and pressure modes, radial and nonradial modes (see, e.g., Christensen-Dalsgaard \\& Houdek 2009).  Here we limit the discussion to single stars (neglecting binary interactions), stellar interiors only (neglecting atmosphere models and winds), and normal phases of evolution (neglecting star formation, supernovae, brown dwarfs, neutron stars, and black holes). ",
        "conclusions": "2D and 3D models are now becoming sophisticated enough to study (and reveal neglected) dynamical processes in stars.  In many cases 3D models generate significantly different results than 2D models.  The 2D and 3D models are being used to generate refinements in parameterized 1D models for longer-timescale evolution calculations.  Updates continue to be made in input and physics for stellar models:  abundances, opacities, EOS, nuclear reaction rates, electron screening, diffusive settling.  Observations from the ground and space are generating profuse data on stellar oscillations that are being used to test the physics and input for stellar interior models.  Many outstanding problems remain, and new ones  revealed by new data on stellar oscillations are emerging.  The research of understanding and simulating the normal evolution and oscillations of single stars is far from complete."
    },
    "1005/1005.0479_arXiv.txt": {
        "abstract": "We re-examine the constraints on the cosmic string tension from Cosmic Microwave Background (CMB) and matter power spectra, and also from limits on a stochastic background of gravitational waves provided by pulsar timing.  We discuss the different  approaches to modeling string evolution and radiation. In particular, we show that the unconnected segment model can describe CMB spectra expected from thin string (Nambu) and field theory (Abelian-Higgs) simulations using the computed values for the correlation length, rms string velocity and small-scale structure relevant to each variety of simulation. Applying the computed spectra in a fit to CMB and SDSS data we find that $G\\mu/c^2< 2.6\\times 10^{-7}$ ($2 \\sigma$) if the Nambu simulations are correct and $G\\mu /c^2< 6.4\\times 10^{-7}$  in the Abelian-Higgs case. The degeneracy between $G\\mu/c^2$ and the power spectrum slope $n_{\\rm S}$ is substantially reduced from previous work. Inclusion of constraints on the baryon density from Big Bang Nucleosynthesis (BBN) imply that $n_{\\rm S} <1$ at around the $4\\sigma$ level for both the Nambu and Abelian-Higgs cases. As a by-product of our results, we find there is ``moderate-to-strong'' Bayesian evidence that the Harrison-Zel'dovich spectrum is excluded (odds ratio of $\\sim 100:1$) by the combination of CMB, SDSS and BBN when compared to the standard 6 parameter fit.  Using the contribution to the gravitational wave background from radiation era loops as a conservative lower bound on the signal for specific values of $G\\mu/c^2$ and loop production size, $\\alpha$, we find that $G\\mu /c^2< 7\\times 10^{-7} $ for $\\alpha c^2/ (\\Gamma\\, G\\mu )\\ll1$ and $ G\\mu/c^2 < 5\\times 10^{-11}/\\alpha$ for $\\alpha c^2/(\\Gamma\\, G\\mu) \\gg1$. ",
        "introduction": "The idea that cosmic strings may have been a source of inhomogeneities in the Universe has been around for some time (see e.g. ref.~\\cite{stringrev} for an overview). Even though it now seems that inflation is the primary source of the inhomogeneities which lead to anisotropies in the Cosmic Microwave Background (CMB), the idea has persisted. The main reason for this is that the amplitude of fluctuations produced by strings formed close to a Grand Unified Theory (GUT) phase transition ($\\sim 10^{16}{\\rm GeV}$) is of the order of observed density inhomogeneities on large scales, making it an attractive probe of physics at energies way above that possible in terrestrial accelerators.  It may seem very non-minimal to have two mechanisms creating inhomogeneities and contrary to the common-sense approach advocated by Occam's Razor. However, there are at least two classes of inflation models, linked to popular ideas in fundamental physics, where strings can be formed and  the two sets of perturbations have similar amplitudes. One is based on Supersymmetric Hybrid Inflation~\\cite{hybrid1,hybrid2,fterm,dterm1,dterm2} and the other is Brane Inflation~\\cite{Dvali:1998pa}. In both cases cosmic strings can be formed during the phase transition at the end of inflation~\\cite{10pera,10perb,10perc}. It was shown in ref.\\cite{Jeannerot:2003qv} that this can happen in a large number of symmetry breaking schemes and it is possible that semi-local strings could form~\\cite{Urrestilla:2004eh,Urrestilla:2007sf} in situations where strings are not forced to exist on topological grounds.  One cannot rule out strings using observations per se, rather one can only constrain the tension (or the mass per unit length) of the strings, and hence the energy scale at which the strings are formed. The tension is usually quantified in terms of the dimensionless ratio $G\\mu/c^2$.  This can be done in many different ways. In this paper we will focus on two: the anisotropies, created by the Kaiser-Stebbins effect~\\cite{Kaiser:1984iv,Fraisse:2007nu} and related pre-recombination processes in the CMB; and the stochastic background of gravitational waves created by the decay of string loops~\\cite{Vachaspati:1984gt}.  There are serious difficulties in setting such limits due to the dynamic range involved in modeling the evolution of a cosmic string network over cosmological scales. This has led to a number of different constraints being published in the literature which are often contradictory -- these will be reviewed in the subsequent sections. The objective of the present paper is two-fold. For both types of observations we will review the important physics and attempt to delineate the key uncertainties. We then present what we believe are the most reliable constraints based on the present knowledge. In the case of the CMB observations we will present two different constraints, based on different assumptions about the evolution of the string network. We note that the constraints presented here will apply only directly to strings where the intercommutation probability, $p$, is unity; that is, they can only be directly applied to normal field theoretic strings and not necessarily to cosmic superstrings~\\cite{superstrings}.   \\begin{figure} \\centering \\mbox{\\resizebox{0.4\\textwidth}{!}{\\includegraphics[angle=0]{TT_str.eps}}} \\mbox{\\resizebox{0.4\\textwidth}{!}{\\includegraphics[angle=0]{cls_comp.eps}}} \\caption{\\label{fig:comp} (Left) Comparison of cosmic string power spectra computed with the Unconnected Segment Model (USM) using  the Nambu -- model A -- (solid line) and AH mimic -- model B -- (dotted line) parameters. Spectra have been normalized to the WMAP value of $C_{10}$, giving $G\\mu  / 10^{-6}c^2=1.18$ for the Nambu model and $1.91$ for the AH mimic. (Right) Comparison of scalar, vector and tensor modes for the USM AH mimic (solid) parameters and the actual AH spectra from simulations (dotted)~\\cite{Bevis:2006mj}. The former has $G\\mu / 10^{-6}c^2=1.91$, and the latter $2.04$. From top-to-bottom at low $\\ell$ the ordering of the curves is the total power, then the anisotropy from vectors, scalars and tensors, respectively.} \\label{fig:spectra} \\end{figure}  \\begin{figure} \\centering \\mbox{\\resizebox{0.7\\textwidth}{!}{\\includegraphics[angle=0]{TT_str_morph_comb.eps}}} \\caption{A montage of plots which show the dependence of the string spectrum on parameters of the USM. In each case the solid and dotted lines are the Nambu and AH mimic models respectively from  Fig.~\\ref{fig:spectra}. In all but the top-left plot we have varied the one parameter keeping all the others the same as in the AH mimic model: (top-left) the dashed line is the USM spectrum with $\\xi=0.13$, $v=0.65c$ and $\\beta=1.9$ but no matter-radiation transition; (top-right) the short dashed line has $\\xi=0.2$ and the long dashed line has $\\xi=0.5$; (bottom-left) the short dashed line has $v=0.0c$ and the long dashed line has $v=0.8c$; (bottom-right) the short dashed line has $\\beta=1.5$ and the long dashed line $\\beta=1.9$.} \\label{fig:morph} \\end{figure} ",
        "conclusions": "There are a wide range limits on $G\\mu/c^2$ presented in the literature, often using identical or at the very least similar data. In this paper we have tried to clarify these issues and present what we believe are current constaints from the CMB and large scale structure data, and also from pulsar timing. We have presented evidence that the string spectra which come from different simulation techniques can all be modelled using the USM with parameters measured in simulations and deduce limits of $G\\mu/c^2< 2.6\\times 10^{-7}$ if the Nambu simulations are correct and $G\\mu/c^2<6.4\\times 10^{-7}$ if the AH simulations are correct. These use CMB+SDSS data and $q_{\\rm str}$ as the parameter. We have also included data from BBN and used $\\log[G\\mu/c^2]$ as the parameter with similar but quantitatively different conclusions.  The limits from pulsar timing could already be much stronger than those from the CMB, for example, if $\\alpha$ is large. However, the uncertainties in interpretation of the pulsar limits on $\\Omega_{\\rm g}h^2$ are much more severe. At this stage, it seems safest to use the low $\\alpha$ limit of $\\Omega_{\\rm g}h^2<8\\times 10^{-7}$,  which is one that is unlikely to have to be taken back.  \\begin{figure} \\centering \\mbox{\\resizebox{0.4\\textwidth}{!}{\\includegraphics[angle=0]{BB_str.eps}}} \\caption{\\label{fig:bb} $B-$mode polarization power spectra from strings, tensors and gravitational lensing. For the string spectra, we use values of $G\\mu/c^2$ corresponding to the $2\\sigma$ upper limit from CMB+SDSS data, that is. $G\\mu / c^2=2.6\\times 10^{-7}$ for the USM Nambu model (solid) and $6.4\\times 10^{-7}$ for the AH case (dotted). We also show the inflationary primordial tensor spectrum with $r=0.1$ at $k=0.05 \\, {\\rm Mpc}^{-1}$ (short dashed) and $r=0.01$ (long dashed). Finally, we show the gravitational lensing spectrum generated from $E-$mode mixing (dot-dash) expected in the inflationary model.} \\label{fig:bmode} \\end{figure}  As a final point we should note that a network of cosmic strings has another possible observational consequence in the CMB, albeit one of lower amplitude: the creation of B-mode polarization~\\cite{Battye:1998js,Bevis:2007qz,Pogosian:2007gi} whose detection to very low values of $G\\mu$ might be possible  if lense cleaning is found to be possible~\\cite{Seljak:2006hi}. This is due to the fact that a substantial fraction of the anisotropies created by strings are vector modes. In Fig.~\\ref{fig:bmode} we have presented the B-mode spectra for the Nambu and AH models with the value of $G\\mu/c^2$ set to their repsective upper limits. These spectra are very different to those predicted in by inflationary models parameterized by the scalar-to-tensor ratio, $r$ which are also included in the figure. In fact, the spectra for $\\ell >30$ are similar to  that produced by gravitational lensing particularly for tte Nambu case. This is because the B-mode polarization is created along the line of sight, albeit on different scales to the gravitational lensing pf the primary E-mode signal. These spectra comprise a white noise portion and a peak corresponding to the dominant scale which is on slightly larger scales in the AH model and, by coincidence, very similar to the lensing for the Nambu model. For the Nambu cas this  similarity will make them difficult to detect unless one can detect the analogue of the reionization bump which peaks around $\\ell\\approx 10$. It is, however, interesting since most inflationary models which produce strings create very low levels of gravitational waves. It has been shown that it is possible to discriminated them in the AH case~\\cite{Urrestilla:2008jv}."
    },
    "1005/1005.4961_arXiv.txt": {
        "abstract": "We report SMARTS, Gemini and \\textit{Swift}-UVOT observations of the optical transient (OT) associated with gamma-ray burst (GRB) 091127, at redshift 0.49, taken between 0.9 hr and 102 days following the \\textit{Swift} trigger.  In our early-time observations, the OT fades in a manner consistent with previously observed GRB afterglows.  However, after 9 days post-burst, the OT is observed to brighten for a period of $\\sim2$ weeks, after which the source resumes fading.  A comparison of this late-time ``bump'' to SN 1998bw (the broad-lined Type Ic supernova associated with GRB 980425), and several other GRB supernovae (SNe), indicates that the most straightforward explanation is that GRB 091127 was accompanied by a contemporaneous SN (SN 2009nz) that peaked at a magnitude of $M_V=-19.0\\pm0.2$.  SN 2009nz is globally similar to other GRB supernovae, but evolves slightly faster than SN 1998bw and reaches a slightly dimmer peak magnitude.  We also analyze the early-time UV-optical-IR spectral energy distribution of the afterglow of GRB 091127 and find that there is little to no reddening in the host galaxy along the line-of-slight to this burst. ",
        "introduction": "The list of gamma-ray bursts (GRBs) that are associated with supernovae (SNe) has grown in the last decade (e.g., GRB 030329/SN 2003dh --- \\citealt{Hjorth+03}, \\citealt{Stanek+03}, \\citealt{Bloom+04}; GRB 031203/SN 2003lw --- \\citealt{Cobb+04}, \\citealt{Gal-Yam+04}, \\citealt{Malesani+04}, \\citealt{Thomsen+04}; GRB 060218/SN 2006aj --- \\citealt{Campana+06}, \\citealt{Ferrero+06}, \\citealt{Mirabal+06}, \\citealt[][and references therein]{Cobb+06a}; GRB 050525a --- \\citealt{DellaValle+06a}; see also \\citealt{Woosley+Bloom06}), indicating that at least some (perhaps most) long-duration GRBs result from the core collapse of massive stars.  However, as evidenced by the failure to detect SNe associated with two nearby long-duration GRBs (GRB 060505 at redshift $z=0.0889$ and GRB 060614 at $z=0.125$: \\citealt{Cobb+06b}, \\citealt{DellaValle+06b}, \\citealt{Fynbo+06}, \\citealt{Gal-Yam+06}, \\citealt{Xu+09}), we do not yet have a complete understanding of the progenitors of long-durations GRBs and the apparent origin of the diversity of their associated SN properties. Long-duration GRB 091127 ($T_{90} = 7.1$\\,s, \\citealt{gcn10191}), presented us with a new opportunity to investigate the GRB-SN connection.  GRB 091127 triggered the \\textit{Swift} Burst Alert Telescope \\citep{Gehrels+04} on 2009 Nov. 27 at 23:25:45 (UT dates are used throughout this paper; \\citealt{gcn10191}).  The prompt emission had a power-law index of $2.05\\pm0.07$ ($-0.4$ to $7.5$ s post-burst) and a fluence of $9.0\\pm0.3\\times10^{-6}$ erg cm$^{-2}$ (15-150 keV) \\citep{gcn10197}. The \\textit{Swift} X-ray Telescope observed the X-ray afterglow with a photon spectral index of $1.98^{+0.15}_{-0.14}$ and absorption column of $9.8^{+3.3}_{-3.1}\\times10^{20}$ cm$^{-2}$ \\citep{gcn10201}.  \\textit{Swift} could not immediately slew to the burst due to Earth-limb constraints, but an optical transient (OT) was quickly identified using the robotic 2-m Liverpool Telescope \\citep{gcn10192}.  \\textit{Swift} Ultraviolet and Optical Telescope (UVOT, \\citealt{Roming+05}) observations, beginning $\\sim 50$ min post-burst, confirmed the existence of this OT at RA$_{\\rm J2000}$ = $02^\\mathrm{h}26^\\mathrm{m}19.89^\\mathrm{s}$, Dec$_{\\rm J2000}$ = $-18^\\circ57^\\prime08\\farcs5$ \\citep{gcn10193}.  Gemini-North Multi-Object Spectrograph (GMOS, \\citealt{Hook+04}) and Very Large Telescope X-shooter spectra of the OT revealed emission features from the underlying host galaxy of the GRB at $z=0.49$ \\citep{gcn10202,gcn10233}.  Because of this relatively low redshift, GRB 091127 was an excellent candidate for follow-up observations in search of a GRB-related SN.  In this {\\it Letter}, we present optical data that show a late-time rebrightening in the OT of GRB 091127.  We interpret this extra component of light as being due to a SN (SN 2009nz, \\citealt{cbet2288}) associated with GRB 091127. The observations, data reduction, and photometry are reported in \\S 2. In \\S 3, we consider each component of light associated with the OT and note the absence of reddening in the burst's host galaxy.  We conclude in \\S 4 with a comparison of past GRB-SNe.  Throughout this paper, we assume the standard cosmological model with $\\Omega_\\Lambda= 0.73$, $\\Omega_m= 0.27$, and a Hubble constant of 71 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": ""
    },
    "1005/1005.1069_arXiv.txt": {
        "abstract": "We present extragalactic number counts and a lower limit estimate for the cosmic infrared background (CIRB) at 15\\,$\\mu$m from {\\it AKARI} ultra deep mapping of the gravitational lensing cluster Abell 2218. These data are the deepest taken by any facility at this wavelength, and uniquely sample the normal galaxy population. We have de-blended our sources, to resolve photometric confusion, and de-lensed our photometry to probe beyond {\\it AKARI}'s blank-field sensitivity. We estimate a de-blended 5\\,$\\sigma$ sensitivity of 28.7\\,$\\mu$Jy. The resulting 15\\,$\\mu$m galaxy number counts are a factor of three fainter than previous results, extending to a depth of $\\sim$\\,0.01\\,mJy and providing a stronger lower limit constraint on the CIRB at 15\\,$\\mu$m of $1.9\\,\\pm$\\,0.5\\,nW\\,m$^{-2}\\,{\\rm sr}^{-1}$. ",
        "introduction": "The cosmic infrared background (CIRB) is dominated by the dusty-emissions from star-forming galaxies, and therefore traces the dust-enshrouded star-formation over the history of the Universe. To interpret the CIRB in terms of galaxy formation and evolution models \\citep[e.g.][]{2004ApJ...600..580G,2005MNRAS.358.1417P}, it is necessary to resolve the monochromatic backgrounds into their individual galaxies. There is a strong correlation between mid-IR and far-IR star-forming galaxies \\citep{2001ApJ...556..562C,2002A&A...384..848E}, therefore galaxies responsible for the CIRB peak $\\sim$140\\,$\\mu$m\\,-\\,200\\,$\\mu$m \\citep{2006A&A...451..417D,2009Natur.458..737D} must also dominate the CIRB at shorter wavelengths, i.e., mid-IR $\\lesssim$\\,60\\,$\\mu$m. Recent results from {\\it Spitzer}, {\\it SCUBA} and {\\it BLAST} data have shown, via stacking analysis of 24\\,$\\mu$m sources at longer wavelengths, that 24\\,$\\mu$m-selected populations account for the bulk of 70\\,$\\mu$m, 160\\,$\\mu$m and 250\\,$\\mu$m backgrounds and also dominate the 350\\,$\\mu$m, 450\\,$\\mu$m and 500\\,$\\mu$m backgrounds \\citep{2006A&A...451..417D,2008MNRAS.386.1907S, 2009Natur.458..737D}. In contrast, at 850\\,$\\mu$m the 24\\,$\\mu$m population only resolves around one quarter of the background \\citep{2008MNRAS.386.1907S}. Previous lower estimates of the CIRB at 15\\,$\\mu$m, from lensed ISOCAM data \\citep[e.g.][hereafter MET03]{2003A&A...407..791M}, have successfully resolved the contribution from galaxies highly luminous in the IR, however the greater depth achieved by these data gives a more representative sample of galaxy populations.\\\\  A major challenge for deep IR observations is fluctuations from confusion noise, which presents a fundamental limit to blank-field surveys \\citep{1974ApJ...188..279C}. Exploiting strong gravitational lensing offers a way to probe beyond the inherent blank-field confusion limit \\citep{1997ApJ...490L...5S}. Where lensing increases the apparent surface area of a background field, the sources within that area are viewed at a lower number density in comparison to the un-lensed situation, and the preservation of surface brightness leads to amplified observed flux densities. Lensing therefore offers a twofold confusion-beating effect. Reduction in the observed area leads to source number-density depletion, and both area correction and flux correction are required to recover the true galaxy-number counts \\citep[][]{1995astro.ph.11150B}. In addition to faint source confusion below the detection limit, a significant proportion of extractions from a confused image may be blends of two or more sources, so photometric de-blending is required \\citep{2006MNRAS.371.1891R}. \\\\ \\\\ In this Letter we present new 15\\,$\\mu$m galaxy number counts and 15\\,$\\mu$m integrated light (IGL$_{15}$) estimate. In Section 2 we summarise the {\\it AKARI} data and data reduction. A data analysis description is given in Section 3, and the results are presented in Section 4 and discussed in Section 5. \\\\ \\\\ Throughout this Letter we assume flat $\\Lambda$CDM cosmology with $\\Omega_{M}$\\,=\\,0.3 and $H_0$\\,=\\,70\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. ",
        "conclusions": "From our de-blended and de-lensed 15\\,$\\mu$m counts we have derived an IGL$_{15}$ estimate of 2.0\\,$\\pm$\\,0.4\\,nW\\,m$^{-2}$\\,sr$^{-1}$, down to $\\sim$\\,0.01\\,mJy. We conclude that, with respect to the P10 model, the {\\it AKARI} 15\\,$\\mu$m data are consistent with having resolved the whole of the predicted IGL$_{15}$. Assuming no radical change between the IR SED of high-redshift galaxies and those resolved at 15\\,$\\mu$m with median redshift of 1.0 \\citep{2002A&A...384..848E}, then the galaxies resolved by these data represent the bulk of galaxies dominating CIRB peak.  Figure 4 suggests that in order to resolve 100\\% of the CIRB$_{15}$, future observations need to probe depths in the region of 1 magnitude fainter than the sensitivity limit achieved by this survey, down to at least {\\it S}$_{15}$\\,=\\,1\\,$\\mu$Jy. The first possible direct measurement constraints of the CIRB$_{15}$ will come from {\\it JWST} or {\\it SPICA} \\citep{2006SSRv..123..485G,2004AdSpR..34..645N}.  15\\,$\\mu$m stacking analysis of {\\it Herschel}/SPIRE and PACS A2218 data, will address how representative 15\\,$\\mu$m selected galaxies are of the galaxy populations responsible for the CIRB at its peak."
    },
    "1005/1005.1296_arXiv.txt": {
        "abstract": "The Ph.D. Thesis of Brian J. Williams, submitted to and accepted by the Graduate School of North Carolina State University. Under the direction of Stephen P. Reynolds and Kazimierz J. Borkowski.\\\\  Interstellar dust grains play a crucial role in the evolution of the galactic interstellar medium (ISM). Despite its importance, however, dust remains poorly understood in terms of its origin, composition, and abundance throughout the universe. Supernova remnants (SNRs) provide a laboratory for studying the evolution of dust grains, as they are one of the only environments in the universe where it is possible to observe grains being both created and destroyed. SNRs exhibit collisionally heated dust, allowing dust to serve as a diagnostic both for grain physics and for the plasma conditions in the SNR. I present theoretical models of collisionally heated dust which calculate grain emission as well as destruction rates. In these models, I incorporate physics such as nonthermal sputtering caused by grain motions through the gas, a more realistic approach to sputtering for small grains, and arbitrary grain compositions porous and composite grains. I apply these models to infrared and X-ray observations of Kepler's supernova and the Cygnus Loop in the galaxy, and SNRs 0509-67.5, 0519-69.0, and 0540-69.3 in the LMC. X-ray observations characterize the hot plasma while IR observations constrain grain properties and destruction rates. Such a multi-wavelength approach is crucial for a complete understanding of gas and dust interaction and evolution. Modeling of both X-ray and IR spectra allows disentangling of parameters such as pre and postshock gas density, as well as swept-up masses of gas and dust, and can provide constraints on the shock compression ratio. Observations also show that the dust-to-gas mass ratio in the ISM is lower by a factor of several than what is inferred by extinction studies of starlight. Future observatories, such as the James Webb Space Telescope and the International X-ray Observatory, will allow testing of models far beyond what is possible now. ",
        "introduction": "Supernova explosions are among the most energetic events in the universe since the Big Bang, releasing more energy ($\\sim 10^{51}-10^{53}$ ergs) than the Sun will release over its entire lifetime. They are the cataclysmic ends of certain types of stars, and are responsible for seeding the universe with the material necessary to form other stars, planets, and life itself. We owe our very existence to generations of stars that lived and died billions of years ago, before the formation of our Sun and solar system. The processes of stellar evolution continue to occur today, with typical galaxies like the Milky Way hosting several supernovae per century, on average. The study of supernovae and the role they play in shaping the evolution of star systems and galaxies is truly an exploration of our own origins. Supernova remnants (SNRs), the expanding clouds of material that remain after the explosion, spread elements over volumes of thousands of cubic light-years, and heat the interstellar medium through fast shock waves generated by the ejecta from the star.  \\subsection{Stellar Evolution}  Any study of SNRs must begin with the processes which cause a star to go supernova. Stars come in all sizes and colors (where the color of a star is related to its temperature), but virtually all stars live out their lives in a similar fashion. They spend the majority of their lives fusing hydrogen in their cores into helium, a process which releases energy. This process is not particularly efficient in stars; nevertheless, the sheer mass of available material to burn means that stars shine at a roughly constant brightness for millions, billions, even tens of billions of years. The life expectancy of a star is a rather sensitive function of its initial mass. This relationship is perhaps counterintuitively inverted, such that the more massive a star is, the shorter its lifetime. (This is due to the fact that massive stars, while having much more fuel to burn, fuse it at a much faster rate than do low-mass stars). Stars that are on the hydrogen burning phase of their lives are said to be on the ``main sequence,'' a reference to the Hertzsprung-Russell diagram seen in Figure 1.1.  This fusion process that takes place in the stellar core is one-half of a balancing act, providing the internal pressure to support the star against the other half, gravity. This cosmic ``tug-of-war'' can only last as long as the star has fuel to burn; once the hydrogen fuel in the core is depleted, gravity will pull the star in on itself unless a new process can provide the necessary counter-balancing force. The fate of a star at this point depends on its mass, with high and low-mass stars following very different paths.  \\subsubsection{Low-Mass Stars}  Stars that begin their main-sequence lives with a mass of less than about 8 solar masses ($M_\\odot$, where $M_\\odot$ is $\\sim 1.99 \\times 10^{33}$ grams) will spend the majority of their lives in the hydrogen burning phase. When their hydrogen runs out, they will swell into red giants, increasing in volume by a factor of $>$ 1000. (The smallest of stars, with masses $<$ 0.5 $M_\\odot$, will not even have enough power to reach this stage). The red giant phase is typically characterized by a helium core surrounded by a hydrogen burning shell. When the core contracts and heats to temperatures above $\\sim 10^{8}$ K, helium burning will begin, fusing helium to carbon via the triple-alpha process. The core of the star that remains typically cannot fuse much beyond carbon and oxygen, and collapses further when the helium supply is used up. The outer layers of the star are ejected, creating the misleadingly named ``planetary nebula.'' Collapse continues until the degenerate pressure of electrons in the plasma is sufficient to balance the gravitational forces, creating a ``white dwarf'' star. A typical white dwarf has a mass of $\\sim 0.6-0.7 M_\\odot$ contained in a volume about the size of the Earth (R $\\sim 6000$ km), which leads to an average density for a white dwarf of $\\sim 10^{6}$ g cm$^{-3}$. White dwarfs are typically very stable (although see section 1.2), and will exist as burned-out remnants of once bright stars indefinitely.  \\subsubsection{High-Mass Stars}  Stars above $\\sim 8$ $M_\\odot$ are hot enough at their cores to fuse hydrogen into helium, helium into carbon, carbon into oxygen, neon, silicon, sulfur and other elements, stepping up the periodic table. Once iron (element 26) is reached, though, it no longer becomes energetically favorable for the fusion process to continue. The iron core collapses in upon itself on timescales of the order of a second, sidestepping electron degeneracy pressure by eliminating electrons, combining them with protons to form neutrons. Once nuclear density ($\\sim 10^{15}$ g cm$^{-3}$) is reached, the degeneracy pressure of neutrons is sufficient to halt the collapse, and the core becomes a proto-neutron star. The remaining infalling material from the core bounces off of this now hard core, ejecting the outer layers of the star in a fantastic explosion known as a ``core-collapse supernova.'' The proto-neutron star forms a neutron star, a stellar remnant of order 1 $M_\\odot$ and R $\\sim 10$ km, with an average density of $\\sim 10^{15}$ g cm$^{-3}$. For highly massive stars, it is possible that even neutron degeneracy pressure cannot halt the collapse of the core, and a stellar-mass black hole is formed.  \\subsection{Supernova Classification}  As in many sciences, observations of events or objects in astronomy often precede theoretical explanations for said events. However, unlike in most disciplines, astronomers normally do not have the means to conduct laboratory tests of observed phenomena. This often leads to a significant time delay between the observation of and the theoretical description of a given event. As a result, the field is riddled with examples of ``historical inaccuracies'' when it comes to naming and classification schemes. A prime example of this is the supernova classification scheme, where astronomers classify supernovae as ``Type I'' or ``Type II'' based solely on the existence of hydrogen lines in their spectra; Type I SNe show no hydrogen lines, Type II do. It was only later realized that vastly different processes can be responsible for this bit of observational data.  \\subsubsection{Type Ia}  The vast majority of stars in the galaxy are $< 8 M_\\odot$, meaning (see Section 1.1.1) that they are destined to end their lives as white dwarfs. However, many stars in the galaxy are also part of binary (or even triple) systems. If a white dwarf is contained in a binary system with a companion star that enters its red giant phase (or has even slightly evolved off the main sequence), and the separation between the stars is sufficiently close, the white dwarf can gravitationally strip matter from its companion. Mass transfer occurs between the two stars, with the white dwarf growing in mass via an accretion disk. White dwarfs can only exist up to the ``Chandrasekhar limit,'' and a white dwarf pushed over this limit ($\\sim 1.4 M_\\odot$) will become unstable, igniting a deflagration, or subsonic burning, of material in the star. This deflagration will lead to a detonation, a supersonic thermonuclear explosion of the entire white dwarf. The entire 1.4 $M_\\odot$ of material is ejected into the surrounding medium at speeds of $\\sim 10,000$ km s$^{-1}$, releasing about 10$^{51}$ ergs of kinetic energy. Type Ia SNe leave behind no compact remnant, and their light curve, i.e. the brightness of the supernova as a function of time, is primarily powered by the radioactive decay chain of nickel-56 to cobalt-56 to iron-56. Their light curves peak a few days after explosion, then slowly fade over the course of a few months. Since all white dwarfs are thought to explode at an identical mass via the same mechanism, their light curves are quite similar in peak brightness, and can be used as a ``standard candle'' to determine extragalactic distances. Type Ia spectra show no hydrogen because white dwarfs are made mostly of carbon and oxygen, nearly all of which is burned to nickel in the explosion.  \\subsubsection{Type II}  Type II supernovae result from the deaths of massive stars, greater than 8 $M_\\odot$, but generally not more than $\\sim 25-30 M_\\odot$. These stars end their lives as core-collapse SNe (CCSNe), ejecting their outer layers (5-25 $M_\\odot$) at a velocity of $\\sim 5,000-10,000$ km s$^{-1}$. Coincidentally, they yield roughly the same amount of kinetic energy ($\\sim 10^{51}$ ergs) as do type Ia SNe, but their overall energetics are much greater. Ninety-nine percent of the energy released in a CCSN is carried off by neutrinos. Type II SNe leave behind a neutron star that is typically $\\sim 1.4 M_\\odot$ with an initial temperature (kT) of several MeV. Type II SNe can be further divided into subclasses based either on the shape of the light curve or the spectrum (e.g. IIP, IIL, IIb, IIn, etc.). They show hydrogen lines in their spectra because the outer atmosphere of the star at the time of explosion still contained a significant amount of hydrogen.  \\subsubsection{Type Ib and Ic}  Type Ib and Ic SNe are also the result of a core-collapse event, differing from type II in that they do not show hydrogen lines in their spectra. The physical reason behind this is that type Ib and Ic SNe result from stars that have shed their outer layers via stellar winds (``Wolf-Rayet'' stars) or gravitational interaction with a companion star prior to explosion. Type Ib explosions result from stars that have lost only their hydrogen envelope; type Ic from stars that have lost both their hydrogen and helium envelopes. Type Ib and Ic SNe are believed to be the result of stars with a progenitor mass of $\\gsim 30 M_\\odot$, although this number could be lower in binary systems. They leave behind neutron stars and stellar-mass black holes. As with type II SNe, most of their energy is carried off in neutrinos, and their light-curves, both in peak brightness and in shape, can vary greatly.  \\subsection{Astrophysical Shocks}  Shock waves, propagating supersonic disturbances, occur commonly in all sorts of astro-environments throughout the universe, where conditions are unlike those found on Earth. Densities in the interstellar medium (ISM) are on the order of a few particles per cubic centimeter, six orders of magnitude less than can be produced in the best laboratory vacuum systems. Sound speeds in the ISM are usually on the order of a few km s$^{-1}$, relatively slow by astrophysical standards. The shock waves generated by a supernova are an excellent example of a strong shock, with shock speeds often being several thousand times the speed of sound in the ISM.  These shock waves compress, sweep, and heat interstellar material. In fact, supernova shock waves are one of the main sources of heating of the interstellar medium, as well as the mechanism for distributing material throughout the universe.  The following is a mathematical description of a shock wave, beginning with the Rankine-Hugoniot Conditions, given by  \\begin{equation} \\rho_1v_1=\\rho_2v_2, \\end{equation}  \\begin{equation} \\rho_1v_1^2+p_1=\\rho_2v_2^2+p_2, \\end{equation}  \\begin{equation} \\frac{1}{2}v_1^2+E_1+\\frac{p_1}{\\rho_1}=\\frac{1}{2}v_2^2+E_2+\\frac{p_2}{\\rho_2}, \\end{equation}  \\noindent where $\\rho$ is the density of the gas, $v$ is the velocity, $p$ is the thermal pressure, and $E$ is the internal energy, given by $dW=dE=-p$d$V$. Subscript 1 denotes pre-shock material (that is, ambient interstellar material that is ahead of the expanding shock wave) and 2 denotes post-shock gas. These equations are written in the frame of reference of the shockwave (see Figure 1.2).  In particular, this means that $v_1$ is not zero (in fact, $v_1 = -v_{shock}$) even though there is no motion in front of the shock in the observer's frame.  Equation (1) is the conservation of mass across the shock, while Equation (2) shows that the sum of the ram pressure, $\\rho v^2$, and thermal pressure $p$ must be equal across the shock. This amounts to a conservation of momentum. Equation (3) is the conservation of energy (kinetic, internal, and thermal).  The volume of the gas is given by $V$, such that d$V=$d$\\frac{1}{\\rho}$. Thus $dE=-p$d$V=-K\\rho^{\\gamma}$d$\\frac{1}{\\rho}$, which can be integrated to obtain $E=\\frac{1}{\\gamma-1}\\frac{p}{\\rho}$. The thermodynamic relation between pressure and density is $p=K\\rho^{\\gamma}$ where $K$ is a constant (although it is not constant across the shock) and $\\gamma$ is the polytropic index of the gas. $\\gamma$ varies depending on the properties of the fluid and is given by: \\\\ $\\gamma=5/3$ for non-relativistic monoatomic gas (as is generally found in the ISM),\\\\ $\\gamma=4/3$ for relativistic monoatomic gas,\\\\ $\\gamma=7/5$ for non-relativistic diatomic gas.\\\\  \\noindent Using this, eq. (3) becomes  \\begin{equation} \\frac{1}{2}v_1^2+\\frac{1}{\\gamma-1}\\frac{p_1}{\\rho_1}=\\frac{1}{2}v_2^2+\\frac{1}{\\gamma-1}\\frac{p_2}{\\rho_2} \\end{equation}  \\noindent Using the Rankine-Hugoniot conditions, it can be shown that  \\begin{equation} \\frac{\\rho_2}{\\rho_1}=\\frac{(\\gamma+1)p_2+(\\gamma-1)p_1}{(\\gamma+1)p_1+(\\gamma-1)p_2}=\\frac{v_1}{v_2}. \\end{equation}  This is the shock jump condition for density, giving the {\\bf compression ratio}, which relates density (and thus fluid velocity) ahead of the shock and behind the shock. In the limit of strong shocks, that is, $p_2 \\gg p_1$, or $M \\gg 1$, where $M$ is the Mach number (defined as $v_{shock}/v_{sound}$), we can neglect the pre-shock pressure so that  \\begin{equation} \\frac{\\rho_2}{\\rho_1}=\\frac{\\gamma+1}{\\gamma-1}=4 \\end{equation}  \\noindent for $\\gamma=5/3$. In the limit of a strong shock, we also have  \\begin{equation} \\frac{p_2}{p_1}=\\frac{2\\rho_1v_1^2}{p_1(\\gamma+1)}, \\end{equation}  \\noindent or  \\begin{equation} p_2=\\frac{2\\rho_1v_1^2}{\\gamma+1}. \\end{equation}  \\noindent The gas is also governed by the ideal gas law such that $p_{2}=n_{2}kT_{2}$, where k is Boltzmann's constant. So  \\begin{equation} p_2=\\frac{2\\rho_1v_1^2}{\\gamma+1}=\\frac{3}{4}\\rho_1v_1^2=n_{2}kT_{2}, \\end{equation}  \\noindent where $n=\\frac{\\rho}{m}$ = total particle number density, and $m$ is the mean mass per particle. Thus,  \\begin{equation} p_2=\\frac{2\\rho_1v_1^2}{\\gamma+1}=\\frac{3}{4}\\rho_1v_1^2=\\frac{\\rho_{2}}{\\mu m_p}kT_{2}, \\end{equation}  \\noindent where $m_p$ is the mass of a proton, equal to $1.67 \\times 10^{-24}$ grams. Further simplifying, we have  \\begin{equation} kT_{2}=\\frac{3}{16} \\mu m_pv_s^2, \\end{equation}  \\noindent where $\\mu m_{p}$ is the mean mass per particle behind the shock (for a fully ionized plasma of cosmic abundances, $\\mu \\sim 0.6$). This is the temperature behind a shock, where $v_{s} = v_1$ is the shock speed. $kT_{2}$ is the ``shock temperature,'' which is the average of proton and electron temperatures. In the absence of heating of electrons at the shock (and assuming no sharing of energy between protons and alpha particles), the initial temperature ratio between protons and electrons, $T_{p}/T_{e}$, is just the ratio of the masses of protons and electrons, $m_{p}/m_{e}$ = 1836. Behind the shock, these temperatures equilibrate to bring down $T_{p}$ and bring up $T_{e}$, but the timescale for equilibration is long, and is a function of the post-shock temperature and density. Observationally, it is known \\citep{ghavamian01} that an inverse correlation exists between shock speed and degree of equilibration of protons and electrons at the shock, with near full equilibration seen in old remnants like the Cygnus Loop ($v_{s}$ $\\sim 300-400$ km s$^{-1}$) and little equilibration seen in younger remnants like Tycho ($v_{s}$ $\\sim 2000$ km s$^{-1}$).  For supernova shock waves of order a few thousand km s$^{-1}$, this shock temperature will be of order 10 million K. Thus, shocked gas will radiate thermal emission at X-ray energies, however, other emission mechanisms also operate in SNRs. Non-thermal emission from synchrotron radiation is often seen near the forward edge of the shock wave in young SNRs, and line emission from ionized elements is common in the shocked ejecta. Shocks in SNRs are ``collisionless,'' in that collisions between particles are extremely rare, meaning that particle interactions are mediated by magnetic fields. This is a good approximation when the Coulomb mean free path is much greater than the gyroradius of the thermal particles.  \\subsection{Supernova Remnants}  The expanding material ejected from the star, rich in heavy elements like oxygen, silicon, and iron, as well as the shock wave which it drives into the ISM is known as a supernova remnant (SNR). Figure 1.3 shows Cassiopeia A, an example of a young ($\\sim 330$ yrs) SNR in our own galaxy. Although Cassiopeia A is known to have resulted from a CCSN, it is generally difficult to tell the type of SN only by looking at the remnant. SNRs remain visible for thousands, often tens or hundreds of thousands of years before dissipating their energy into the ISM. The life of a SNR can be thought of as consisting of four phases.  \\subsubsection{Free Expansion Phase}  Immediately following the explosion, the ejecta from the supernova race out into the ISM at speeds of $\\sim 10,000$ km s$^{-1}$, driving a strong shock at the leading edge. Since the ejected mass is significantly greater than the mass of the rarefied medium it encounters, there is no appreciable slowing of the ejecta by the ambient medium. How long this phase lasts depends on both the ejected mass of the SN and the density of the CSM in front of the shock, and can be anywhere from a few days to years.  \\subsubsection{Reverse-Shock Phase}  As the shock continues to expand and sweep material it encounters, the accumulated mass becomes non-negligible, and gradually causes the shock to slow. The ejecta behind the shock, however, are still traveling at a free-expansion velocity, and slam into the decelerating material ahead of it. This causes a reverse shock to form. Initially, this ``reverse'' shock moves inward only in the Lagrangian shock frame, and still moves outward in the observer's frame. The reverse shock, however, eventually ``turns around'' and moves inward in the frame of reference of the observer. It typically takes of order hundreds of years for this transition to occur. When the ejecta are in free-expansion cooling is almost entirely adiabatic. While this adiabatic cooling is effective in lowering the temperature of the ejecta, energy is still conserved for the system because the energy is not radiated away. Upon encountering the reverse shock the ejecta are heated, like the ISM at the forward shock, to very high temperatures, thus radiating strongly in X-rays. This radiative cooling is still relatively inefficient, and most of the energy of the SNR+ISM system is conserved. The duration of the reverse shock phase can last tens to thousands of years.  \\subsubsection{Sedov-Taylor Phase}  Once the mass swept-up by the forward shock greatly exceeds the ejecta mass, the remnant enters the Sedov-Taylor phase (often known simply as the Sedov phase). The reverse shock has propagated all the way back through the ejecta and dissipated, and the remnant can be described by a self-similar solution (Sedov 1959). The similarity variable can be derived by dimensional analysis, and is given by  \\begin{equation} \\xi = R(\\rho/Et^2)^{1/5} \\end{equation}  \\noindent where $R$ is the distance the blast wave has traveled from the supernova, $E$ is the explosion energy of the initial event, $\\rho$ is the density of the ISM, and $t$ is time. $\\xi$ is dimensionless, and equation 12 can be used to show that the distance traveled by the blast wave (for $\\gamma = 5/3$) as a function of $E$, $\\rho$, and $t$, is given by  \\begin{equation} R \\propto (\\frac{E}{\\rho})^\\frac{1}{5} t^\\frac{2}{5}. \\end{equation}  \\noindent It can be immediately seen from this that the shock velocity is given by  \\begin{equation} V_s=\\frac{dR}{dt}=\\frac{2R}{5t}. \\end{equation}  \\noindent The Sedov phase lasts for thousands to tens of thousands of years after the explosion.  \\subsubsection{Radiative Phase}  The shock continues to sweep up material and decelerate, eventually reaching a point where the forward shock speed is only a few hundred km s$^{-1}$. At this point, the temperature of the post-shock gas drops below 10$^{6}$ K, and radiative cooling of the gas becomes important. Cooling of the gas is a runaway process, as the more it cools, the more the cooling rate increases. As the gas temperature drops further, the material once again becomes visible in optical radiation. The forward shock is driven mostly by momentum conservation at this point, and eventually will turn sub-sonic and dissipate into the ISM.  \\subsection{Radiation Mechanisms in SNRs}  SNRs radiate throughout the electromagnetic spectrum. In radio waves, the emission is entirely non-thermal in origin, resulting from synchrotron radiation from relativistic electrons spiraling around magnetic fields. Synchrotron emission is characterized by a featureless power-law spectrum, where the radio flux, $S_{\\nu}$ is given by $S_{\\nu} \\propto \\nu^{-\\alpha}$, where $\\nu$ is the frequency and $\\alpha$ is the spectral index, which depends on the energy distribution of the electron population.  The primary focus of this work is on infrared (IR) emission from SNRs, the physical basis for which is described in Chapter 2. Briefly, IR emission is dominated by thermal continuum from warm dust grains, heated by collisions with the hot ions and electrons in the post-shock region. Remnants in the radiative phase can show IR line emission as well, from low to moderately ionized states of abundant heavy elements like O, Ne, S, Si, Ar and Fe.  At optical wavelengths, radiation prior to the radiative phase comes primarily from hydrogen Balmer lines (transitions from $n \\ge 3 \\to 2$), such as H$\\alpha$, $\\lambda$ = 656.3 nm, and H$\\beta$, $\\lambda$ = 486.1 nm. This requires the presence of neutral hydrogen ahead of the shock, which is much more easily attained in the case of a type Ia SN, since CC SNe generally ionize the surrounding medium, either with ionizing radiation from the progenitor, or a flash of ultraviolet (UV) radiation at the moment of explosion. Charge exchange between slow neutral atoms and fast protons behind the shock produces fast-moving neutral atoms, generating a broad H$\\alpha$ line, with a narrow component arising from stationary neutral atoms in the post-shock medium. During the radiative phase, strong optical lines are seen from a variety of atomic species, most strongly from H$\\alpha$ and singly-ionized sulfur ([S II]). Ultraviolet (UV) emission from SNRs is also produced, generally from higher ionization states than in optical.  Soft X-rays (0.1-2 keV) are generally thermal in origin, and in SNRs are often dominated by line emission from highly ionized elements. Typically, elements that emit X-ray line emission have been stripped of all but one or two electrons, making them ``hydrogen-like'' or ``helium-like.'' As with optical and infrared lines, downward transitions of electrons to lower energy levels causes emission of a photon whose energy is equal to the transition energy between the electron's bound states. For elements that still contain multiple electrons, the transitions between energy states become more complicated, and generally generate numerous lines which are smeared together by current X-ray spectroscopic technology. Continuum emission, detailed below, is also observed at these energies.  Hard X-rays (2-50 keV) in SNRs can be either thermal or non-thermal in origin. Thermal X-rays are dominated by continuum emission, although lines from K-shell electron transitions do exist beyond 2 keV for elements such as Si, S, Ar, and Fe. Thermal continuum seen in X-ray spectra is generally thermal Bremsstrahlung (also known as ``free-free'' emission), which occurs when electrons and protons interact. This process causes the electron to slow down, and energy conservation requires emission of a photon to account for the lost kinetic energy of the particle. Free-bound emission (or radiative recombination) can also take place when a proton or ion captures a free electron, emitting in the process a photon whose energy depends on both the free kinetic energy of the electron and the orbital to which it is captured.  Non-thermal emission in SNRs arises from synchrotron emission identical to that seen in radio waves, but from much more energetic electrons. The maximum photon energy in keV of an electron with energy $E$ is given by  \\begin{equation} h \\nu = 1.93 (E/100 TeV)^{2} (B/10\\mu G) keV. \\end{equation}  \\noindent In order to produce synchrotron emission in the 2-10 keV range, electrons with energies of 100-200 TeV are required. Because this is well beyond the particle thermal energies for even a fast shock, another process must accelerate particles to high energies in remnants where non-thermal emission is conclusively identified. The origin of these high-energy cosmic-ray electrons is discussed in the next section.  At gamma-ray energies, emission can be produced by one of three processes; two of which are leptonic in origin, one of which is hadronic. Bremsstrahlung, both thermal and non-thermal in origin, can account for photons of all energies, up to TeV emission. Inverse-Compton scattering from relativistic electrons off of cosmic microwave background or far-infrared photons can upscatter the photons to very high energies. The only known hadronic source of gamma-ray emission is the decay of neutral pions, or $\\pi^{0}$ particles, into two gamma-ray photons. This process occurs 98.7\\% of the time in $\\pi^{0}$ decays. The $\\pi^{0}$ particles themselves are produced in collisions between cosmic-ray protons and thermal protons, as well as protons and alpha particles in the pre-shock gas. The minimum pion energy required to produce a gamma-ray of energy $E_{\\gamma}$ is given by  \\begin{equation} E_{min}(\\pi) = E_{\\gamma} + (m_{\\pi}^{2}c^{4}/4E_{\\gamma}) \\end{equation}  \\noindent To produce high energy gamma rays ($> 1$ GeV), the last term on the right becomes negligible, and the minimum pion energy needed is roughly equal to the gamma-ray energy observed.  It is likely that all of these processes play a role in the gamma-ray emission observed from SNRs. For a thorough review of SNRs at high energies, see Reynolds (2008).  \\subsection{Cosmic-Ray Acceleration in SNRs}  Cosmic-rays are highly energetic particles, typically protons, alpha particles, and nuclei of heavier elements, with a small percentage of the population consisting of electrons, streaming through space at relativistic speeds. They can either be detected directly (at lower energies), or indirectly through interactions with atoms in the Earth's upper atmosphere (at high energies). Upon the collision of a cosmic-ray with our atmosphere, a shower of particles (mostly containing pions) is produced. These pions decay further into electrons, positrons, muons, neutrinos, and photons, and can be detected from ground-based Cherenkov telescopes. Such telescopes can even reconstruct the events to determine the location in the sky from which the cosmic-ray came. Unfortunately, this location is not indicative of the original source of the particle. Since cosmic-rays are charged particles, they gyrate around the magnetic field lines of our Galaxy, and are essentially randomized by the time they reach Earth. This leads to a mystery: from where do cosmic-rays come?  \\subsubsection{Cosmic-Ray Sources}  The fact that synchrotron emission is observed in radio waves for every Galactic SNR known shows that, at the very least, electrons are efficiently accelerated to energies of a few GeV. If electrons are accelerated, protons and ions should be accelerated as well. This, unfortunately, is a difficult thing to observationally verify, because the synchrotron radiation from relativistic protons is orders of magnitude weaker than from electrons spiraling around a magnetic field. Gamma-ray production via the decay of $\\pi^{0}$ particles, discussed in the previous section, requires that protons at GeV energies exist. Unambiguous detection of this hadronic gamma-ray signal in SNRs would provide the observational confirmation that such acceleration of ions is taking place. Searches for this signal are currently underway.  Detection of non-thermal synchrotron emission at X-ray energies is a clear indication that the shock is accelerating electrons beyond TeV energies. If the shocks are equally as efficient at accelerating protons, this could account for the galactic cosmic-ray spectrum to energies up to $\\sim 10^{15}$ eV. Cosmic-rays at much higher energies have been detected, but it is difficult to produce them in large quantities in SNR shocks. It is likely that these ultra-high energy particles have an extra-galactic origin.  If supernova shock waves are efficiently accelerating cosmic rays, then the equations detailed in Section 1.3 are no longer valid, since escaping cosmic rays can rob the shock of energy. This leads to a higher compression ratio (r $\\equiv \\rho_{2}/\\rho_{1}$), and a lower post-shock temperature. Even if no cosmic rays escape, the compression ratio can still be increased if relativistic particles dominate, since r $\\to$ 7 as $\\gamma \\to \\frac{4}{3}$.  \\subsection{Summary}  Supernovae represent the end of a star's life, but in the process of dying, elements that will go on to form future generations of stars and planets are spread throughout the galaxy. The universe is nearly 14 billion years old, old enough that every cubic centimeter of a galaxy like the Milky Way has been overrun numerous times by shock waves produced by SNe. They represent one of the main feedback mechanisms in the evolution of a galaxy, shaping and recycling products in the ISM.  SNRs are visible at all wavelengths of the electromagnetic spectrum, though the physical processes responsible for emission at various wavelengths differ. Nonetheless, these processes are often connected, and a complete understanding of the dynamics of the remnant requires connecting the physics behind these various faces of the remnant.  \\newpage  \\begin{figure} \\figurenum{1.1} \\includegraphics[width=15cm]{fig1-1_small.eps} \\caption{Hertzsprung-Russell Diagram, showing the temperatures of stars vs. luminosity. Taken from {http://www.le.ac.uk/ph/faulkes/web/stars/o\\_st\\_overview.html}. \\label{hrdiagram} } \\end{figure}  \\newpage  \\begin{figure} \\figurenum{1.2} \\includegraphics[width=17cm]{fig1-2_small.eps} \\caption{Shockwave diagram, in the frame of reference of the shock. Arrows represent the velocities of particles upstream and downstream of the shock. In this frame, upstream material (at rest in the observer's frame) races in towards the shock. The shock then randomizes the velocities of the particles, and gives them a bulk velocity downstream of 1/4 their initial velocity (for a strong shock with $\\gamma = 5/3$). \\label{shockframe} } \\end{figure}  \\newpage  \\begin{figure} \\figurenum{1.3} \\includegraphics[width=15cm]{fig1-3_small.eps} \\caption{SNR Cassiopeia A, seen in X-rays from the {\\it Chandra X-ray Observatory}. Image from {http://chandra.harvard.edu/photo/2006/casa/} \\label{casa} } \\end{figure}   \\newpage \\clearpage ",
        "conclusions": "Shock waves generated by supernova explosions play a crucial role in the evolution of the ISM. By heating gas and dust, spreading new elements throughout the cosmos, and reshaping the dust and gas-phase abundances, SNRs are one of the dominant feedback mechanisms for galaxies, and can mark not only the endpoint, but also the beginning of stellar evolution. By studying SNRs, we can learn about both the conditions and dynamics of the remnants themselves and their impact on their surrounding medium. This work has focused on the need for a multi-wavelength approach to studying SNRs, with emphasis on the IR portion of the spectrum, where we see emission from warm dust grains heated via collisions with hot protons and electrons in the post-shock gas. IR studies of SNRs can shed light on one of the primary mysteries of ISM physics: How much dust is present in the ISM, and where was it formed?  \\subsection{Dust-to-Gas Mass Ratio}  SNRs studied in this document are in both the Milky Way galaxy and the LMC, of both core-collapse and type Ia origin, and span a range in ages from young (e.g. 0509-67.5, Kepler) to middle aged (e.g. DEM L71, N132D) to old (e.g. N49B, Cygnus Loop). The shock ages, pre-shock densities, post-shock densities, and amount of material swept vary over orders of magnitude for various objects. However, every remnant studied shares one curious property: the dust-to-gas mass ratio inferred for the ambient ISM is lower than what is inferred from general absorption studies along lines of sight to bright UV sources in the galaxy and Magellanic Clouds.  Dust masses are determined from photometric and/or spectroscopic observations of IR emission from SNRs. When a model spectrum is fit to data, the overall normalization of the model is provided by only two parameters: the radiating dust mass and the distance to the object. For LMC remnants, the distance is well known to be $\\sim 50$ kpc. For galactic remnants there is a bit more uncertainty in the distance, but for Kepler and the Cygnus Loop these uncertainties are likely less than 25\\%. Since SNR shock waves destroy dust in addition to lighting it up, I calculate the amount of sputtering that has taken place for a given set of shock parameters and use this to infer the initial mass in dust warmed by the forward shock. The dust mass arrived at via this method is the amount of dust present in the {\\it ambient ISM}.  Gas masses swept up by the forward shock can be determined in several ways. The X-ray emissivity of a hot plasma is proportional to the product of electron density and the total mass in gas present. If the density can be determined, then the swept gas mass can be calculated. Alternatively, if one has a measure of the pre-shock density in the ambient ISM from optical/UV observations, one can multiply by the volume to get the total gas mass overtaken by the blast wave, either for the entire remnant or for a small portion of the shock.  Dust-to-gas ratios derived in this work are between a factor of 3 and 20 lower than what is expected for the ISM of both the Milky Way and LMC. Obviously, the dust masses derived from model fits are model dependent, but unless major new physics is involved, varying parameters in current models cannot produce enough variation in results to accommodate conventional dust-to-gas ratios. {\\em Even simple analytic expressions that relate the observed IR luminosity to the total amount of radiating dust reproduce dust masses within 50\\% of those derived using the more sophisticated modeling techniques described here.} In recent work on the Cygnus Loop \\citep{sankrit10}, we took an alternative approach, determining the path length through the emitting region for a small section in the NE corner of the Loop. By comparing model fluxes with observed values, we found that a low dust-to-gas ratio is required to derive a path length consistent with FUSE and ROSAT data. The ratio determined in this fashion was 0.38\\%, a factor of 2-3 lower than the standard galactic value.  It is possible that some dust remains cold even behind the shock, and would escape detection in the short-wavelength MIPS and IRS bandpasses on {\\it Spitzer}. I will discuss this possibility more below as it relates to the total amount of dust produced in SNe. Regarding the dust-to-gas mass ratio issue, most of the derived values reported in this work include a detection from the forward shock in the 70 $\\mu$m {\\it Spitzer} band. For a significant amount of dust to be present, it would have to be quite cold to escape detection at 70 $\\mu$m. For dust in the post-shock environment to be cold, it must be contained in dense clumps. Such densities would imply radiative shocks, which are easily detectable via optical emission. It is difficult to imagine a scenario where dust could exist at such cold temperatures while immersed in the hot plasma of a non-radiative shock.  \\subsection{Ejecta Dust in SNRs}  IR and sub-millimeter observations of high redshift galaxies imply large amounts of dust present in the early universe \\citep{nozawa03}. The source of this dust remains a mystery, and it is believed that SNe must play a major role in synthesizing dust grains. \\citet{nozawa03} considered Population III SNe and concluded that CC SNe in the early universe would need to convert 2-5\\% of the mass of the progenitor into dust grains to account for observations. In their models, a 20 $\\msun$ progenitor would produce $\\sim 0.6$ $\\msun$ of dust, while a 30 $\\msun$ progenitor would yield $\\sim 1.3$ $\\msun$. For pair-instability SNe with progenitor masses $> 140 \\msun$, 20-80 $\\msun$ of dust may be produced. In more recent work, \\citet{cherchneff10} find that their models of dust formation in ejecta produce about a factor of 5 less dust than in Nozawa et al., with a 20 $\\msun$ progenitor star only producing $\\sim 0.1-0.15$ $\\msun$ of new dust. Since we cannot make direct observations of SNe at that distance, we are left to examine more local SNe to determine the amount of dust formed in the ejecta.  {\\em I find no evidence for a large amount of ejecta dust in any SNR studied thus far.} Although in some cases the possibility of a large quantity of cold dust cannot be eliminated, we can conclusively rule out a significant amount of warm dust in the ejecta in every object. In Kepler's SNR, the 400-year old remnant of a type Ia explosion, we find an upper limit of 0.1 $\\msun$ based on both 70 and 160 $\\mu$m {\\it Spitzer} data. Further observations at longer IR wavelengths will be needed to confirm this result. In 0540-69.3, the young remnant of a core-collapse SN, we find only $\\sim 10^{-3} \\msun$ of dust in the inner ejecta, although it is possible that more could be unshocked.  While it is certainly possible that cold dust could exist in the ejecta and remain undetected by {\\it Spitzer's} instruments if it were unshocked, it is difficult to conceive of a scenario where dust passes through the reverse shock and remains cold. This again would require a radiative shock, which would be easily detectable. Examination of Figure 4.1 shows no visual evidence of dust morphologically associated with X-ray emitting ejecta for remnants of CC SNe. To measure this quantitatively, I use upper limits at 70 $\\mu$m for several of the middle-aged remnants of CC SNe reported in Chapter 4. To get an upper limit on the amount of ejecta dust, it is necessary to choose a temperature for the dust. \\citet{rho08} report observations of 0.02-0.05 $\\msun$ of freshly-formed ejecta dust in Cas A. The temperatures of the various ``cold'' dust components observed range from 40-150 K. I adopt a temperature of 50 K for dust in the ejecta that has been through the reverse shock, and report upper limits to the dust mass that could be present. I consider the entire volume of the remnant interior to the inner edge of the forward shock.  For N132D, which is by far the brightest SNR in the IR in the sample considered, the upper limit on dust at 50 K is $\\sim 0.15 \\msun$. This is comparable to the amount predicted from CCSNe by \\citet{cherchneff10}. \\citet{france09} used UV data from N132D to derive a progenitor mass of 50 $\\msun$. \\citet{nozawa03} predict a dust formation mass of 2-5\\% of the progenitor mass from CC SNe, leading to a mass of 1-2.5 $\\msun$ from a progenitor of 50 $\\msun$. In order for this mass of dust to be present in N132D, and be below upper limits at 70 and 160 $\\mu$m from {\\it Spitzer}, it would have to be below $\\sim 35$ K. For N49B and 0453-68.5, both of which are fainter by nearly an order of magnitude, the upper limit on dust masses are $\\sim 0.03$ and $\\sim 0.015 \\msun$, respectively. These are only intended to be rough estimates, as the composition and temperature of ejecta dust is unknown. Nonetheless, it appears unlikely that current SNe are the prodigious dust producers that Pop. III SNe may have once been. It is possible that large amounts of dust are initially {\\it formed} in the ejecta from supernovae, but do not survive passage through the reverse shock, but dust created and destroyed in this fashion would not increase the overall dust abundance in a galaxy. The data from {\\it Spitzer} do now allow us to determine this, and observations with the longer wavelength detectors of {\\it Herschel}, with its better resolution and sensitivity than {\\it Spitzer} at wavelengths beyond 100 $\\mu$m, are necessary to confirm these results, as well as examine the possibility of large amounts of cold dust previously undetected.  \\subsection{Future Work}  Typical models for ISM grains assume spherical particles with solid material filling fractions of unity all made of one material. This is probably unrealistic, as real grains almost certainly contain some degree of {\\bf porosity (i.e. part of their volume is taken up by vacuum)} (Okamoto et al. 1994). It is also likely that grain formation mechanisms allow for {\\bf composite grains, i.e. grains made up of multiple types of materials} (Shen et al. 2009). Grains that are porous, composite, or both have properties that are different from their ``compact'' counterparts. Their optical constants are different, and must be approximated with either an effective medium theory (EMT), where different optical constants are added together by various means (see Bohren and Huffman, 1983 for a review of various EMTs), or something like the multilayer sphere approach developed by Voshchinnikov \\& Mathis (1999), or the discrete dipole approximation developed by Draine and Flatau (1994). In work done on this front thus far, I have found an EMT to be the most straightforward way to approximate the dielectric constants for arbitrary materials. I have already developed code that is capable of calculating the optical constants and absorption coeffecients for grains of any arbitrary combination of astronomical silicates (Draine \\& Lee 1984), graphite, amorphous carbon, and vacuum. This code can be extended, if necessary, to include other grain ingredients.  One also needs to modify {\\bf rates of energy deposition and sputtering} for porous and composite grains to properly calculate emission from these grains. The field of porous and composite grains is a relatively new and wide open field of astrophysics, and opportunities to test models have in the past been few and far between. SNR properties inferred from a porous grain model can be significantly different than those from a compact model, as shown in Chapter 8. The potential to learn about the nature of dust in the ISM {\\bf and} the nature of SNRs is great. The case of the young SNR 0509-67.5 (0509) is one example where porous grains may show a better fit to data than do compact grains.  For supernovae and young SNRs, dust emission is often seen at {\\bf short wavelengths}, such as in IRAC energy bands. All attempts to obtain a self-consistent model to explain this emission in relation to that in the mid-IR have been unsuccessful (Blair et al. 2007). The problem is related to modeling the {\\bf stochastic nature of small grains}, which are primarily responsible for short wavelength emission. For sufficiently small grains with small heat capacities, grain temperatures fluctuate between subsequent collisions, and grains emit most of their energy at the highest temperatures. These small grains are particularly sensitive to the {\\bf energy deposition rates of impinging particles}, which regulate both the heating and destruction of grains. These rates are typically assumed to be simple functions of the properties of the grain and the incoming particle (Dwek \\& Werner 1981), but laboratory data have shown that such an approach is not valid under extreme conditions, such as those seen in young SNRs. Since small grains are also {\\em by far} the most numerous grains in the ISM, any model of dust must have more realistic models for the physics involved. For young SNRs, like 0509, ion temperatures are sufficiently high compared to electrons (Warren \\& Hughes 2004, Ghavamian et al. 2007) that {\\bf heating of grains by ions} is non-negligible. I have already begun exploring this, but as with the porous grain calculations, I believe I have only scratched the surface on this front. It will be absolutely crucial to understand short-wavelength IR emission before {\\it JWST} is launched, as it will only observe from $\\sim 1-20$ $\\mu$m, and thus will only be sensitive to emission from small grains. I will develop and implement more physically realistic approaches to energy deposition rates, particularly for small grains.  As grains are sputtered, their constituent elements are {\\bf liberated back into the gaseous phase}. Although dust, by mass, makes up a tiny fraction of the ISM, the depletion of heavy elements onto dust grains can be significant. Any attempt to model X-ray spectra as a function of post-shock ionization timescale should take these increasing abundances for elements like Fe, Mg, and Si into account (Vancura et al. 1994). Our code does allow us to follow this liberation, and I would implement this effect in X-ray spectral modeling of SNRs, while exploring ways to make such information available to the community as a whole.  Work remains to be done on both the theoretical and observational side to better understand SNRs in both the IR and X-ray bands. Models need to be further tested against archival data and results from laboratory astrophysics, particularly in regard to heating and sputtering of porous grains. With the cold phase of {\\it Spitzer} now over, its usefulness in this particular field is minimal, but huge datasets of observations of the galactic plane and Small and Large Magellanic Clouds were completed and archived before the cryogen was exhausted. New observatories are already operating or will be in the next few years. The diffraction-limited optics of {\\it Herschel} (a factor of 4 better than {\\it Spitzer}), combined with its increased sensitivity at long wavelengths, make it an ideal candidate to differentiate between various grain models and find any cold dust that is present in SNRs. At shorter wavelengths, {\\it SOFIA}, while not as sensitive as {\\it Spitzer}, will provide much clearer views of bright galactic objects, allowing tighter constraints on dust destruction behind the shock. The launch of {\\it JWST} in the middle of this decade will truly usher in a new era of IR astronomy, with a resolution nearly an order of magnitude better than {\\it Spitzer}. Since {\\it JWST} will focus on the 1-25 $\\mu$m range of the spectrum, much work needs to be done in the modeling of emission from small grains, including those which thermally fluctuate between collisions. And although we are still more than a decade away from {\\it IXO}, both {\\it Chandra} and {\\it XMM-Newton} are still operating, and the amount of archival data from both of these observatories is vast. The last 6 years of {\\it Spitzer} observations have revolutionized our understanding of the IR window on the universe, and we have only scratched the surface of what the next 10 years will bring.  \\newpage \\clearpage"
    },
    "1005/1005.4685_arXiv.txt": {
        "abstract": "A rapid timing analysis of \\vlt/ULTRACAM (optical) and \\rxte\\ (X-ray) observations of the Galactic black hole binary GX~339--4 in the low/hard, post-outburst state of June 2007 is presented. The optical light curves in the $r'$, $g'$ and $u'$ filters show slow ($\\sim$ 20 s) quasi-periodic variability. Upon this is superposed fast flaring activity on times approaching the best time resolution probed ($\\sim$ 50 ms in $r'$ and $g'$) and with maximum strengths of more than twice the local mean. Power spectral analysis over $\\sim$ 0.004--10 Hz is presented, and shows that although the average optical variability amplitude is lower than that in X-rays, the peak variability power emerges at a higher Fourier frequency in the optical. Energetically, we measure a large optical vs. X-ray flux ratio, higher than that seen on previous occasions when the source was fully jet-dominated. Such a large ratio cannot be easily explained with a disc alone. Studying the optical:X-ray cross-spectrum in Fourier space shows a markedly different behaviour above and below $\\sim$0.2 Hz. The peak of the coherence function above this threshold is associated with a short optical time lag with respect to X-rays, also seen as the dominant feature in the time-domain cross-correlation at $\\approx$150 ms. The rms energy spectrum of these fast variations is best described by distinct physical components over the optical and X-ray regimes, and also suggests a maximal irradiated disc fraction of 20 per cent around 5000 \\AA. If the constant time delay is due to propagation of fluctuations to (or within) the jet, this is the clearest optical evidence to date of the location of this component. The low-frequency quasi-periodic oscillation is seen in the optical but not in X-rays, and is associated with a low coherence. Evidence of reprocessing emerges at the lowest Fourier frequencies, with optical lags at $\\sim$10 s and strong coherence in the blue $u'$ filter. Consistent with this, simultaneous optical spectroscopy also shows the Bowen fluorescence blend, though its emission location is unclear. But canonical disc reprocessing cannot dominate the optical power easily, nor explain the fast variability. ",
        "introduction": "The radiative output of active X-ray binaries (XRBs) peaks at X-ray wavelengths, hence their name. A characteristic property of this radiation is its strong and aperiodic variability over a range of time scales \\citep[e.g. ][]{vanderklis89}. This is a manifestation of rich underlying structure in the accretion environments of these sources on a range of physical scales. X-ray timing studies can thus provide key information for our understanding of volatile XRB environments, and X-ray timing is now a mature field in this respect \\citep[e.g. ][]{hasingervanderklis89,vanderklis95,remillardmcclintock06}.  At lower energies, e.g. in the optical and infrared, the picture is thought to be complex. For many low mass XRBs, the donor star contribution can be ignored during active accretion states. Optical emission from the accretion disc arises predominantly from its outer, cool portions with temperatures much below $\\sim$10$^6$ K. Consequently, large variations of this component should occur only on a relatively-slow viscous timescale. Any faster variations may be the result of high-energy photons being reprocessed on the outer disc. In fact, such X-ray heating is thought to be the primary engine generating the bulk of the ultraviolet and optical emission seen from XRBs \\citep{vanparadijsmcclintock94}. Weaker, but highly significant reprocessing on the companion star is also an important effect which can be used to determine the physical parameters of the binary system \\citep[e.g. ][]{horne85, marsh88, obrien02}.  Observations are now also beginning to provide evidence of fast variability (on $\\sim$1~s timescales or less) at low energies which is not associated with reprocessing, and recent works attribute much of this to non-thermal contributions from the jet and corona. Some observations of fast low energy variations in several sources have actually existed for many years \\citep[e.g.][]{motch82, uemura02_v404cyg, fender97}. In several cases, simultaneous low and high energy observations argue against a reprocessing origin. This was best demonstrated in \\xtej1118, for which rapid optical and X-ray timing observations revealed complex flux variations, including an anti-correlation between optical and X-rays -- exactly the opposite of that predicted by reprocessing scenarios \\citep[][]{kanbach01, spruitkanbach02}. Many works have successfully modelled the optical power and broad-band spectrum as a result of (cyclo-)synchrotron emission, though the location and characteristics of the emitting plasma remain a matter of debate \\citep[e.g. ][ and more]{merloni00, esin01, markoff05, yuan05}. Self-consistently incorporating the timing properties in this framework has proven to be trickier. An energy reservoir model feeding a common jet and corona system developed by \\citet{malzac04} can reproduce both the shape and the strength of the observed timing correlations very well. But physical modeling of the reservoir itself and of the impulses of energy release governing the observed variability remains somewhat ad hoc, being modelled simply as a linear superposition of fluctuating stochastic shots. \\citet{g09_rmsflux} discusses a simple modification of this phenomenological model to reproduce another observable (the \\lq rms--flux\\rq\\ relation), by making the light curves a non-linear coupling of shots, rather than an additive superposition.  Other sources with intriguing multi-wavelength timing correlations include GRS~1915+105, studied by \\citet{eikenberry98} and \\citet{arai09}, among others. And in a recent study, \\citet[][ hereafter Paper I]{g08} and \\citet{durant08} observed two other Galactic BH candidates in the X-ray low/hard state, \\gx339\\ and \\swiftj1753, simultaneously in X-rays with the {\\sl Rossi X-ray Timing Explorer} (\\rxte) and in the optical with ULTRACAM mounted on the Very Large Telescope (VLT). Cross-correlating the light curves in the two wavebands revealed complex variability patterns in both systems. \\citet{durant09} have discussed the second source in detail, suggesting that optical cyclotron anti-correlates with an X-ray (presumably inverse Compton dominated) response (e.g. \\citealt{fabian82}). In the present paper\\footnote{Based on observations obtained during ESO programs 079.D-0535 and 279.D-5021.}, we present the full spectral and timing properties of \\gx339, supplementing the cross-correlation function presented in Paper I.  GX~339--4 (4U~1658--48) is one of the best-studied Galactic black hole (BH) candidates over a wide range of wavelengths. The source is a classic X-ray transient. It is usually found in the active, low/hard flux state, but is known to undergo dramatic changes on month--year timespans, showing a broad range of X-ray states from \\lq quiescence\\rq, to the \\lq very high\\rq\\ state \\citep[e.g. ][]{zdziarski04, dunn08}. Such changes are accompanied by strong modulations in its radio \\citep{corbel00}, and optical \\citep{makishima86} properties. The source has proven to be key for the discovery and elucidation of several important aspects of XRB behaviour. It was one of the first to show clear evidence of the existence of an inner accretion disc in the X-ray low/hard state \\citep{miller06}. It is the prototype source for defining an important relationship between the radio and X-ray fluxes of XRBs \\citep[e.g. ][]{corbel00, gallo03}.  It was also one of the first sources found to show very fast optical flickering \\citep{motch82}. Follow-up rapid simultaneous optical and X-ray studies uncovered evidence of a slow anti-correlated response between the two bands, but such studies remained scarce \\citep{motch83, makishima86}. In Paper I, we presented a new correlation feature in the faint low/hard state, with a sharp, sub-second optical response in a red ($r'$) filter, delayed with respect to X-rays. As was discussed in that paper, and elaborated upon herein, simple linear reprocessing transfer functions cannot explain this feature. We now present our full multi-filter variability dataset, and relate this with the simultaneous X-ray timing properties over a range of timescales. The optical and X-ray power spectra and correlations in time and Fourier space are analysed. Supplementary optical spectroscopy with the FORS2/\\vlt\\ instrument is also presented. The observations are discussed in the context of multiple varying components with some (as yet not fully understood) underlying interactions. Our broad-band analysis provides first simultaneous constraints on the physical characteristics, emitted powers {\\em and} relative locations of these components within the accretion environment.  The donor star of this low mass XRB system is very faint \\citep{shahbaz01}, which complicates measurement of physical parameters of the binary, but means that the optical flux is completely dominated by the accretion flow, which is our main interest herein. We assume a compact object mass of 6~\\Msun, consistent with the best determination of the mass function of 5.8\\p0.5 \\Msun \\citep{hynes03_gx339}, and a distance of 8 kpc (a range of 6--15 kpc has been suggested in the literature; e.g. \\citealt{hynes04, zdziarski04}). The most-widely accepted value of the binary orbital period is 1.75 days, based on outburst radial velocity variations of the optical Bowen emission blend reprocessed on the secondary star \\citep{hynes03_gx339}, and we adopt this value here. ",
        "conclusions": "Our wide optical/X-ray spectral and timing simultaneous coverage gives new insight on the processes occurring on fast timescales in the inner accretion regions of this important source. Fast optical timing is beginning to probe interactions of the jet, corona and disc. This provides good starting material for detailed modelling of the accretion environment. But many questions remain to be answered, including the origin of the optical QPO, the correlation of the observed lags with the various variability components found in the PSDs, the origin of the optical emission lines, and the underlying driver of the fast fluctuations and observed power in \\gx339, and also other sources. The present inferences are based on relatively short datasets, with the best simultaneous dataset having a total duration of less than one hour. Longer follow-up observations in various states of the sources should enhance the constraints presented herein. Indeed, several systems are now seen to show interesting variability characteristics on fast times, and such work ought to continue with new instruments being mounted on large telescopes and offering such modes on an increasing basis: e.g. FORS/VLT high time resolution mode, SALT/SALTICAM and Ultraspec, among others. Further insight could come from improving the time resolution of the observations, or through detailed dynamical cross-correlation analyses utilising each segment of the light curves separately (cf. Durant et al. 2010, in prep.). Coordination with the infrared offers another excellent opportunity, as this may be able to probe optically-thick jet emission. Indeed, \\citet{casella10} report the detection of fast infrared flaring in GX~339--4 with the infrared also delayed by $\\sim$100 ms with respect to X-rays. Finally, instruments such as the LAXPC on board \\astrosat\\ will provide increased hard X-ray sensitivity as compared to \\rxte/PCA\\ \\citep{astrosat} and thus also help to break some of the model degeneracies discussed herein."
    },
    "1005/1005.3770_arXiv.txt": {
        "abstract": "The color-magnitude relation of early-type galaxies differs slightly but significantly from a pure power-law, curving downwards at low and upwards at large luminosities ($M_r>-20.5$ and $M_r<-22.5$, respectively).  This remains true of the color-size relation, and is even more apparent with stellar mass ($M_* < 3\\times10^{10} M_\\odot$ and $M_* > 2\\times 10^{11}M_\\odot$, respectively). The upwards curvature at the massive end does not appear to be due to stellar population effects. In contrast, the color-$\\sigma$ relation is well-described by a single power law. Since major dry mergers change neither the colors nor $\\sigma$, but they do change masses and sizes, the clear features observed in the scaling relations with $M_*$, but not with $\\sigma > 150 \\, {\\rm km~s}^{-1}$,  suggest that $M_* > 2\\times 10^{11}M_\\odot$ is the scale above which major dry mergers dominate the assembly history.   We discuss three models of the merger histories since $z\\sim 1$ which are compatible with our measurements.  In all three models, dry mergers are responsible for the flattening of the color-$M_*$ relation at $M_* > 3\\times10^{10} M_\\odot$ -- wet mergers only matter at smaller masses. At $M_* > 2\\times 10^{11}M_\\odot$, the merger histories in one model are dominated by major rather than minor dry mergers. In another, although both major and minor mergers occur at the high mass end, the minor mergers contribute primarily to the formation of the ICL, rather than to the stellar mass growth of the central massive galaxy. This model attributes the fact that $\\alpha < 1$ in the scaling $M_*\\propto M_{\\rm dyn}^\\alpha$, to the formation of the ICL. A final model assumes that the bluest objects today were assembled by minor dry mergers of the bluest (early-type) objects at high redshift, whereas the reddest objects were assembled by a mix of major and minor dry mergers.  In this model, the scatter of the color-magnitude relation should increase with redshift, and the dependence on environment should also be more pronounced at higher redshift:  more clustered objects should be redder. Similar measurements of these relations at high redshift will provide further valuable constraints on the mass scale at which major dry mergers dominate the assembly history. ",
        "introduction": "The colors of early-type galaxies are tightly correlated with their luminosities (Sandage \\& Visvanathan 1978).  The mean relation is well-described by a single-power law whose slope evolves little out to $z\\sim 1$ (e.g. Kodama et al. 1998; Mei et al. 2009). This, and the small scatter around the mean relation, are thought to imply that the stellar populations in these objects are old (e.g. Bower et al. 1992; Bernardi et al. 2003b,c), although the total stellar metallicity, the $\\alpha$-elements-to-iron abundance ratio, and light-weighted age, all increase along the relation (e.g. Bernardi et al. 2006; Gallazzi et al. 2006). Data sets are now large enough that significant departures from simple power laws can be detected:  the mean color-magnitude relation appears to be steeper at faint luminosities (e.g. Baldry et al. 2004; Graham 2008; Skelton et al. 2009). This change in slope is thought to indicate that the mechanism by which the stars were assembled into a single object is different at low luminosities than at higher ones.  However, different morphological types define different color-magnitude relations.  Since the mix of morphological types is a strong function of luminosity, it is possible that the observed curvature is really due to morphology, rather than to a change in formation histories at fixed morphology. Unfortunately, it is difficult to select large samples of a given morphological type that are pure.  In what follows, we compare the color-magnitude relation obtained from a number of different ways of defining an early-type sample.  We argue that while the steepening of the relation at faint luminosities may be affected by morphological effects, it appears to be present even in relatively pure samples of ellipticals -- this may arise from the fact that dwarf and giant ellipticals are known to be different in other ways. However, we also show that, at the very highest luminosities, $M_r < -22.5$, the relation steepens again. In a companion paper, Bernardi et al. (2010b) show that this steepening occurs on the same scale where the size-luminosity and velocity dispersion-luminosity relations steepen and flatten, respectively (Bernardi et al. 2007).  In addition, the trend for axis-ratio and color-gradient to increase with luminosity, reverse on this scale (Bernardi et al. 2008; Roche et al. 2010). When expressed in terms of stellar mass, the relevant scale is $M_* = 2\\times 10^{11}M_\\odot$.  Section~\\ref{sample} describes the SDSS sample, and a number of ways for selecting early-types from it. Section~\\ref{bent} presents the associated color-magnitude relations, and shows that the trends we see are even more pronounced if we replace luminosity with stellar mass. It also shows that, in contrast, the color-$\\sigma$ relation is well-described by a single power law over essentially the entire range of $\\sigma$. Section~\\ref{sec:ageZ} shows that the curvature is not due to stellar population effects. Section~\\ref{sec|theory} compares our empirical results with simple models.  While these toy models are not intended to provide a precise quantification of the color evolution, they provide a useful framework within which to discuss our measurements. A final section summarizes our findings, and discusses what they suggest about how the formation and assembly of early-type galaxies depend on mass and redshift.  Appendix~\\ref{2gs} describes a way of selecting early-types which exploits the fact that galaxy properties are approximately bimodal; Appendix~\\ref{fukugita} contrasts this with selection based on eyeball classifications of morphology. A number of tests of systematics -- robustness to changes in the scale on which rest-frame color is measured or inferred (color gradients and $k+e$ corrections) -- are described in Appendix~\\ref{systematics}. A final Appendix provides details of the expected changes to galaxy sizes and velocity dispersions if galaxy mergers occur along parabolic orbits and conserve mass and energy.  Where necessary we assume a flat background geometry that is dominated at the present time by a cosmological constant $\\Lambda_0=1-\\Omega_0$, where $\\Omega_0 = 0.3$ is the background density in units of the critical density, with Hubble constant $H_0 = 70$~km~s$^{-1}$Mpc$^{-1}$. ",
        "conclusions": "\\label{discuss} Our study of the color-magnitude and color-$M_*$ scaling relations has revealed interesting trends: one at $M_*\\sim 3\\times 10^{10}M_\\odot$ which had been noticed before (Kauffmann et al. 2003; Skelton et al. 2009), and another at high luminosities ($M_r\\le -22.5$) and masses ($\\log_{10}(M_*/M_\\odot)\\ge 11.3$), which is new to our work. These trends are qualitatively independent of exactly how the early-type sample is selected.  In most cases, the (weak) dependence on precisely how the sample was selected can be traced to contamination of the red-sequence by edge-on spirals. In a related paper, Bernardi et al. (2010b) show that a number of other scaling relations also indicate that these luminosity and mass scales are special.  The red sequence is considerably straighter and narrower than the blue (Figure~\\ref{gmrM}).  However, it is not a simple power law:  it is shallower between $-20.5>M_r>-22.5$ than at either the fainter or brighter ends (Figure~\\ref{gmrEtypes} and Table~\\ref{gmrMtable}). This curvature is not due to contamination by later morphological types at the faint end (Figure~\\ref{gmrMorph}).  It also does not depend on whether one uses {\\tt Petrosian} or {\\tt model} colors (Figure~\\ref{gmrMPM}); although color-gradients mean that the scale on which the color is defined does lead to small quantitative differences.  The curvature is also robust to (reasonable changes in) the choice of $k$-correction provided one properly accounts for evolution (Figure~\\ref{gmrEvolM}).  Unless care is taken to account for it, this curvature may be confused with evolution in magnitude limited surveys (discussion following Figure~\\ref{gmrEvolM}). All these properties of the color-magnitude relation are also true of the color-stellar mass relation (Figures~\\ref{gmrEtypes}, \\ref{gmrM}, \\ref{gmrMs} and Tables~\\ref{gmrMtable} and~\\ref{gmrMsredblue}), and the color-$R_e$ relation (Figure~\\ref{gmrSize}).  The curvature towards redder colors at the brightest ($M_r\\le -22.5$), most massive ($\\log_{10}(M_*/M_\\odot)\\ge 11.3$) end is evident at fixed age and metallicity, suggesting that it is not driven by stellar population effects (Figure~\\ref{fig|AgeMet}). In contrast, the color-$\\sigma$ relation shows no curvature at high $\\sigma > 150 \\, {\\rm km~s}^{-1}$ (Figure~\\ref{gmrSigma}). The fact that there is no feature at the largest $\\sigma$, despite clear features in the scalings with $M_*$, has strong implications for models of the assembly histories of massive galaxies.      Skelton et al. (2009) have argued that the change from a steeper slope at low luminosities to a shallower one at $M_r<-20.5$ is due to a change in formation histories.  They associate the shallower slope with recent major dry mergers which are expected to increase the luminosity and stellar mass without changing the color significantly. Since such mergers are expected to leave the velocity dispersion unchanged, that fact that there is no curvature in the color-$\\sigma$ relation (Figure~\\ref{gmrSigma}) seems in striking agreement with the dry major merger hypothesis.  In addition, dry major mergers are expected to increase the size in proportion to the mass, and we do see some flattening in the color-$R_e$ relation (Figure~\\ref{gmrSize}). However, if the flattening at intermediate luminosities (and stellar masses, and sizes), with no curvature in the color-$\\sigma$ relation is indeed due to major dry mergers, then it seems difficult for such a scenario to explain the steepening at even higher luminosities ($M_r<-22.5$ or $\\log (M_*/M_\\odot) > 11.3$), even though these are precisely the objects for which the dry merger hypothesis is most commonly invoked.  Therefore, we discussed three models that are compatible with our measurements:  one in which major mergers dominate the mass growth at $M_* > 2\\times 10^{11}M_\\odot$ (Figure~\\ref{fig|Mssigma1}), another in which mergers are both major and minor, but the minor mergers at these largest masses contribute to the intracluster light (Figure~\\ref{fig|Mssigma2}), and a third in which the reddest most massive objects today, which happen to also be the oldest, formed from major and minor mergers of the oldest, reddest objects in the past (Figure~\\ref{fig|Mssigma3}), whereas the bluest objects formed from minor (but not major) mergers of blue objects.  Observations of the scatter and environmental dependence of the color-$M_*$ relation at $z\\sim 1$, and of the color-$R_e$ relation at intermediate sizes, will discriminate between these models.  (The color-$\\sigma$ relation is useful too; we are assuming it is harder to measure at high $z$.)  Such tests, e.g., using the thickness of the red sequence to constrain the formation histories of early-type galaxies, must be done with care. This is because although samples defined by cuts in concentration alone may provide a reliable estimate of the mean shape of the red sequence, they provide a bad estimate of the thickness (see Figures~\\ref{gmrEtypes} and~\\ref{gmrM}).  In particular, the red sequence in such samples is thicker at fainter luminosities, because of contamination by edge-on galaxies.  Appendix~\\ref{2gs} shows that double-Gaussian fits to the bimodal color-magnitude relation, while purely statistical, provide a simple way of correcting approximately for this contamination. For example, at intermediate luminosities (i.e., around $L_*$), this procedure correctly assigns the reddest objects to the blue cloud, rather than to the red sequence (Figures~\\ref{histgmr} and~\\ref{histgmrMs}). Appendix~\\ref{fukugita} shows that such objects tend to be edge-on spirals, and can be a significant source of contamination if one simply defines the red sequence by a straight color cut (Figure~\\ref{gmrMorph}).  While they can also easily be removed by a cut on axis ratio (e.g. require $b/a \\ge 0.6$), cutting on concentration index instead does not remove these objects (compare Figures~\\ref{histgmrab} and~\\ref{histgmrci}).  In contrast to samples defined by color or concentration, the width of the red-sequence defined by double-Gaussian fits is independent of luminosity.  In our dataset, we find this width to be $0.033$~mags (Table~\\ref{gmrMredblue}). Since measurement errors are of order 0.02~mags, the intrinsic width may be more like $0.026$~mags.  Our results suggest that, to obtain results which are less likely to be biased by contamination, it is this width which should be compared with the analogous quantity in higher redshift samples. On the other hand, we found that the double-Gaussian decomposition was not able to account for about 10\\% of the objects at luminosities below $M_r<-20.5$; these objects tended to populate the green valley between the red and blue sequences.  So, if the double-Gaussian fits are to be used at higher redshift, one must first check that such objects are not much more common than they are at $z\\sim 0.1$.  Our models assume that massive objects have experienced major mergers since $z\\sim 1$, meaning that the total stellar mass in early-types with $M_*>2\\times 10^{11}M_\\odot$ today must have been smaller by about a factor of 2 at $z\\sim 1$.  It is not clear that this is consistent with current constraints.  E.g., although Faber et al. (2007) claim that the number density of early-types has increased by a factor of at least two since $z\\sim 1$, and Matsuoka \\& Kawara (2010) argue that the number density of objects with $M_*>2\\times 10^{11}M_\\odot$ has increased by an order of magnitude since $z\\sim 1$, Brown et al. (2007), Wake et al. (2008) and Cool et al. (2008) claim that the mass growth since $z\\sim 0.7$, for objects with $M_*>10^{11}M_\\odot$, must have been less than 50\\%. Eliche-Moral et al. (2010) argue that some of the discrepancy between these two claims is due to the difference between how the samples were defined.  However, most of these constraints are based on parametrizations of the $z\\sim 0$ stellar mass function which may have underestimated the true abundance at $M_*>2\\times 10^{11}M_\\odot$ by 50\\% (see Bernardi et al. 2010a and references therein).  If the true local abundance is indeed larger, then major mergers may be {\\em required} to reconcile the $z\\sim 1$ counts with those at $z\\sim 0$.  Finally, it is interesting to ask how BCGs, which are amongst the most massive objects in the local universe, fit into this picture? Although we do not show them explicitly here, they define a similar color-$M_*$ relation (and other relations as those shown in Figure~1 of Bernardi et al. 2010b) for $\\log (M_*/M_\\odot) > 11.3$. However, there are some important differences: compared to non-BCGs of similar mass or luminosity, their colors are slightly redder (Figure~\\ref{fig|Mssigma3}, and Roche et al. 2010), they have smaller color gradients (Roche et al. 2010), and slightly larger sizes (Bernardi 2009). Whereas the first two suggest merger histories dominated by major mergers, consistent with their large masses, the fact that their sizes are larger suggests more size growth than is usually associated with major mergers.  This suggests that although major mergers erased their color gradients at some higher redshift, minor mergers have puffed up their sizes, decreased their velocity dispersions further and contributed to the formation of ICL at lower redshift (Bernardi 2009)."
    },
    "1005/1005.4827_arXiv.txt": {
        "abstract": "We report on a 70ks {\\xmm} ToO observation of the newly discovered accreting millisecond pulsar, {\\igr}. Pulsations at 244.8339512(1) Hz are observed throughout the outburst with an RMS pulsed fraction of 14.4(3)$\\%$. Pulsations have been used to derive a precise solution for the P$_{orb}=12487.51(2)$s binary system.  The measured mass function indicates a main sequence companion star with a mass between 0.15 and 0.44 M$_{\\odot}$.  The {\\xmm} 0.5--11 keV spectrum of {\\igr} can be modelled by at least three components, which we interpret, from the softest to the hardest, as multicoloured disc emission, thermal emission from the neutron star surface and thermal Comptonization emission. Spectral fit of the {\\xmm} data and of the {\\it RXTE} data, taken in a simultaneous temporal window, well constrain the Comptonization parameters: the electron temperature, kT$_e$=51$^{+6}_{-4}$ keV, is rather high, while the optical depth ($\\tau$=1.34$^{+0.03}_{-0.06}$) is moderate.  The energy dependence of the pulsed fraction supports the interpretation of the cooler thermal component as coming from the accretion disc, and indicates that the Comptonizing plasma surrounds the hot spots on the neutron star surface, which in turn provide the seed photons.  Signatures of reflection, such as a broadened iron K$\\alpha$ emission line and a Compton hump at $\\sim 30$ keV, are also detected.  We derive from the smearing of the reflection component an inner disc radius of $\\simgt$ 40 km for a 1.4 M$_{\\odot}$ neutron star, and an inclination between 38$^{\\circ}$ and 68$^{\\circ}$.  {\\xmm} also observed two type-I X-ray bursts, whose fluence and recurrence time suggest that the bursts are ignited in a nearly pure helium environment.  No photospheric radius expansion is observed, thus leading to an upper limit on the distance to the source of 10 kpc. A lower limit of 6.5 kpc can be also set if it is assumed that emission during the decaying part of the burst involves the whole neutron star surface.  Pulsations are observed during the burst decay with an amplitude similar to the persistent emission. They are also compatible with being phase locked to pre-burst pulsations, suggesting that the location on the neutron star surface where they are formed does not change much during bursts. ",
        "introduction": "The discovery of 244.8 Hz pulsations in the X-ray emission of the newly discovered transient source, {\\igr} \\citep{Mrk09a}, has brought to  thirteen the number of Accreting Millisecond Pulsars (hereafter AMSP) discovered so far.  Such high spin frequencies are believed to ensue from a long phase ($\\simgt 10^8$ yr) of accretion of mass and angular momentum onto the neutron star (NS), according to the so-called recycling scenario \\citep[][]{BhtvdH91}.  In this paper we mainly focus on the spectral properties of this source. The broadband continuum of AMSP has been invariably found to be dominated by a power-law like hard emission with a cut off at an energy between 30 and 60 keV .  This emission is interpreted as Comptonization of soft photons in a hot plasma \\citep[see][and references therein, for a review of the spectral properties of AMSP]{Pou06}.  At energies of $\\sim$1 keV two soft components are also generally found (\\citealt{GrlPou05}, GP05 hereinafter, \\citealt{Pap09,Pat09}). The cooler one is attributed to the accretion disc emission, while the hotter is interpreted as thermal emission of the hotspots on the NS surface. In this scenario, the hotspot thermal emission constitutes the bulk of soft photons that Compton down-scatter the hot plasma.  The X-ray spectrum of {\\igr} is found consistent with this phenomenological description and the outlined physical interpretation fits well the observed spectral properties.   \\begin{figure} \\includegraphics[angle=270,width=\\columnwidth]{F1.eps} \\caption{2--60 keV light curve of the {\\igr} outburst as observed by the PCU2 of the PCA aboard {\\rxte}. Circles and crosses refer to observations pointed towards {\\igr} (Obsid P94041) and {\\xte} (Obsid P94042), respectively. As the two PCA fields of view greatly overlap, the detected pulse frequency is used to discriminate from which source the emission observed by the PCA was actually coming from (see text for details). To which source the emission is attributed is indicated at the top of the figure. The time at which the {\\xmm} observation was performed is also shown.} \\label{fig:lc} \\end{figure}  A couple of AMSP have also showed the typical clues of disc reflection, a Compton hump at $\\sim30$ keV \\citep{GrlDonBar02} and a K$\\alpha$ emission iron line \\citep{Pap09,Cac09,Pat09}. In particular the line observed from {\\saxj} is compatible with the typically broadened shape caused by the relativistic Keplerian motion of the inner rings of the accretion disc, immersed in the deep gravitational well of the compact object \\citep{Fab89}. The observation of similar broadened lines from AMSPs has a peculiar importance in the context of NS accretors.  It may constrain in fact where the magnetosphere breaks off the accretion disc flow and lifts off matter to the accretion columns.  The spectral and timing capabilities of {\\xmm} already proved fundamental to investigate the X-ray emission and pulse properties of AMSPs especially at low energies, as well as to detect iron emission lines at a high statistics. To these ends, we have obtained a 70ks Target of Opportunity (ToO) observation of {\\igr}, performed roughly one week after the outburst onset. We report on the spin and orbital properties of the source in Sec. \\ref{sec:pulse}, and on its spectrum in Sec.\\ref{sec:xmmspectrum}.  We also show how it is possible to take advantage of a simultaneous Rossi X-ray Timing Explorer ({\\it RXTE}) observation, in order to constrain its high energy emission (Sec. \\ref{sec:rxtexmm}).  Also two type I X-ray bursts were caught by {\\xmm}, making this observation particularly rich (Sec.\\ref{sec:burst}).  We observed burst oscillations during the bursts, at the same frequency of the non-burst emission.  It is of key importance to observe burst oscillations from sources, the rotational state of which is known at a great accuracy. The comparison with the properties of the non-burst pulsations may in fact assess the physical mechanism that produces oscillations during bursts.  These have already been observed to be strongly phase locked to the non-burst oscillations, possibly suggesting that the similarities in their formation are more profound than expected \\citep{Wts08}.  \\section[]{Observations and data analysis}  The X-ray transient {\\igr} has been discovered by {\\it INTEGRAL} on 2009 September 21.140, during its Galactic Bulge monitoring \\citep{Bald09}. At that time also the {\\it RXTE} Proportional Counter Array (PCA) detected a rising X-ray flux activity from the Galactic Bulge. This emission was tentatively attributed to two known nearby transient sources, {\\xte}, an AMSP spinning at 435 Hz \\citep{Mark02}, and GRS 1747--312. {\\rxte} therefore pointed at the position of {\\xte} to discover that the emission instead originated by a source spinning at $\\sim$245 Hz. This made {\\igr} the twelfth AMSP discovered \\citep{Mrk09a}. A follow up {\\it Swift} ToO observation constrained the source position at a few arcsec accuracy, obtaining a value $\\sim$20 arcmin away from {\\xte}, thus confirming the discovery of a new source \\citep{Bozzo09}. A {\\it Chandra} observation further refined the source position \\citep{Now09}, giving the most accurate available estimate as RA=17h~51'~08.66'', DEC=-~30$^{\\circ}$~57'~ 41.0'' (1 sigma error of 0.6 arcsec), which is the one we consider in the following.  \\subsection{\\xmm}   \\begin{figure} \\includegraphics[angle=270.0,width=\\columnwidth]{F2.eps} \\caption{Pulse phase time delays induced by the orbital motion as a function of the mean orbital longitude. Longitude 0$^{\\circ}$ refers to the passage of the pulsar at the ascending node of the orbit. Residuals with respect to the best fitting orbit model are also over-plotted at a magnified scale (rightmost y-axis). The reduced chi squared of the model is 1.15 over 684 degrees of freedom.  } \\label{fig:orbit} \\end{figure}  {\\xmm} observed {\\igr} as a ToO observation for 70ks starting on 2009 September 21.037, $\\sim$ 9 days after the source discovery. The {\\xmm} observation is indicated by a horizontal bar in Fig.\\ref{fig:lc}, where the {\\it RXTE} PCA light curve of the outburst is plotted.  The EPIC-pn camera operated in timing mode to allow the temporal resolution (30 $\\mu$s) needed to study the millisecond variability of the source, and to avoid spectral deformations due to pile-up. Also EPIC-MOS2 camera observed in timing mode, while EPIC-MOS1 operated in Small Window to provide an image of the source.  No evident background due to soft proton flares was recorded during the observation, and the full exposure window is retained for scientific purposes. {\\xmm} detected two type I X-ray bursts during its pointing.  When analysing the {\\it persistent} emission\\footnote{Even if the source is an X-ray transient, we refer to {\\it persistent} emission as the main body of the outburst emission, and to {\\it burst} emission when the source exhibits thermonuclear flashes, so called type I X-ray bursts.}, we discarded 20s prior and 110s after the burst onset. This choice is fairly conservative as the e-folding factor of the burst decay is $\\tau\\simeq11$s (see Sec.\\ref{sec:burst}).  {\\xmm} data have been extracted and reduced using SAS v.9.0.  A concatenated and calibrated event file has been created with the task epproc, specifying the Chandra coordinates of the source.  The EPIC-pn spectrum has been extracted from a 13 pixels wide region around the source position (RAWX=36), equivalent to 53.3 arcsec (which should encircle more than 90 per cent of the source counts up to 9 keV\\footnote{See {\\xmm} Users handbook, issue 2.6, available at http:// xmm.esac.esa.int/external/xmm\\_user\\_support.})  A similar stripe centred on RAWX=10 has been used to estimate the background. The background subtracted 1.4--11 keV\\footnote{see Sec.\\ref{sec:xmmspectrum} for details about the energy interval chosen for the purposes of a spectral analysis.} {\\it persistent} count rate recorded by the EPIC-pn is 43.9 c/s on the average.  Pile up is therefore not an issue for EPIC-pn data.  Spectral channels have been re-binned to have at least three channels per resolution element and 25 counts per channel.  EPIC-MOS1 data have been extracted using the task emproc to provide a 101''$\\times$101'' image of the source. Using the task edetect\\_chain, we estimate the source position as RA=17h 51m 08.55s, DEC=-30$^{\\circ}$ 57' 41.7'', with an uncertainty of 4.3 arcsec.  This is consistent with the more accurate Chandra determination. EPIC-MOS 2 data have a significantly lower statistics than that of EPIC-pn and are therefore discarded from the spectral analysis.  Reflection Grating Spectrometers (RGS) operated in standard spectroscopy mode and data have been reduced using the task rgsproc. We consider only first order spectra, re-binned in order to have at least 25 counts per channel. The average {\\it persistent} count rate is 0.4 and 0.5 c/s for RGS1 and RGS2, respectively.  All spectral fits were performed using XSPEC v.12.5.1  \\subsection{\\rxte} \\label{sec:rxte}  Similarly to the other AMSPs, only one third of the X-ray flux of {\\igr} is emitted in the energy range covered by {\\xmm}. In order to better constrain the broadband X-ray emission, we take advantage of two {\\it RXTE} observations performed simultaneously to the {\\xmm} pointing, and amounting to an exposure of 8.8ks (Obs. 94041-01-02-08 and Obs. 94041-01-02-18).  {\\rxte} observations have also been used to extract a light curve of the whole outburst. Starting on 2009 September 13.849, {\\rxte} observed {\\igr} for $\\sim$ 24.5 d, achieving a total exposure of 611 ks ($\\simeq$30\\% of the whole outburst; ObsId P94041). The 2--60 keV light curve extracted from the PCU2 of the PCA \\citep{Brd93,Jah06} data is plotted in Fig.\\ref{fig:lc}. Background has been subtracted using the {\\it faint} model (pca\\_bkgd\\_cmfaintl7\\_eMv20051128.mdl), regardless during the outburst the source crosses the threshold of 40 c/s/PCU above which a {\\it bright} model should be used. Such a choice was made to avoid the unphysical flux discontinuity that would take place when switching from one model to the other. In Fig.\\ref{fig:lc}, filled circles represent observations pointed to the {\\igr} position (ObsId P94041), while crosses refer to pointings in the direction of the nearby source, {\\xte} (ObsId P94042). These two sources are 19.67 arcmin away, and there is an overlap of nearly 80\\% between the respective PCA pointings (the FWHM of the PCA collimator is $\\sim$ 1$^{\\circ}$). {\\it RXTE} observations are subject to the contribution of both sources, no matter which the instrument is actually pointed to. The detected spin frequency is therefore used to discriminate which source emits the observed flux.  As during the first observation reported in Fig.\\ref{fig:lc} the 245 Hz periodicity is detected, it can be attributed to {\\igr} even if the instrument was pointing towards {\\xte}. The overlap between {\\rxte} observations of these two sources becomes even more evident when looking at the late part of the outburst. Around 2009 Oct 8 (Day 26 according to the scale used in Fig.\\ref{fig:lc}, where time is reported in days since 2009 Sep 12.0), the X-ray flux shows an increase of more than 70\\% with respect to the previous emission.  Most importantly, 435 Hz pulsations are detected, clearly indicating a renewed activity from {\\xte}, as it has also been confirmed by a narrow pointed {\\it Swift} XRT observation \\citep{Mrk09b}. Observations performed between 2009 Oct 8 and Oct 10 can be therefore safely attributed to a dim outburst of {\\xte}, similar to those already displayed in 2005 March \\citep{Grb05} and 2007 April \\citep{Fal07}.  The fact that {\\igr} is in the Galactic Bulge not only forces to deal with a crowded field, but also with the Galactic Ridge emission. It is in fact easy to see from Fig.\\ref{fig:lc} that the PCU2 count-rate stays at a level of $\\sim$ 6.5 c/s even after the {\\xte} activity episode is over. Emission at the same flux level was observed also at the end of the 2002 outburst of {\\xte} (GP05) and owes to the Galactic ridge. This emission  introduces an additional background to all the PCA observations. This is clearly indicated by the prominent 6.6 keV iron line that appears in PCA spectra when this additional background is not accounted for \\citep{Mrk09a}. As we show in Sec.\\ref{sec:rxtexmm}, to subtract this background is of key importance to ensure agreement between the EPIC-pn and RXTE-PCA spectral data.  To estimate a reliable spectrum of this additional background we therefore summed up all {\\rxte} data from 2009 Oct 11.400 to 22.745 (for an exposure of 87.8 ks). In the following we refer to this additional background as the Galactic ridge emission.  In order to extract a PCA spectrum, we have considered only data taken by the PCU2, that is the only always switched on during the {\\rxte} observations overlapping with the {\\xmm} pointing. Data processing has also been restricted to its top xenon layer, as it is the less affected by the instrumental background. Only photons in the 3--50 keV band have been retained as this band is best calibrated. Data taken by the High Energy X-ray Timing Experiment \\citep[HEXTE,][]{Rth98} in the 35--200 keV band has been extracted considering the cluster B, which was the only one to perform rocking to estimate the background, at the time the observation took place.  \\begin{table} \\caption{Best fitting orbital and spin parameters of {\\igr} as observed by {\\xmm}. Errors in parentheses are evaluated at 1$\\sigma$ confidence level, while upper limits are given at a 3$\\sigma$ level.} \\label{table:spin} \\centering \\renewcommand{\\footnoterule}{}  % \\begin{tabular}{lr} \\hline a $\\sin{i}$/c (lt-s) & 0.275196(4)   \\\\ P$_{orb}$ (s) & 12487.51(2)\\\\ T$^*$ (MJD) & 55094.9695351(7)\\\\ e & $<8\\times10^{-5}$\\\\ f(M$_1$,M$_2$,i) (M$_{\\odot}$) & $1.07086(5)\\times10^{-3}$\\\\ \\hline $\\nu$ (Hz) & 244.8339512(1) \\\\ $\\dot{\\nu}$ (Hz/s) & $<3\\times10^{-11}$\\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "\\subsection{The X-ray pulsations}  X-ray coherent pulsations at a frequency $\\nu=244.8339512(1)$ Hz are observed throughout the observation performed by {\\xmm}. No spin frequency evolution is detected at a 3 $\\sigma$ upper limit of $1\\times10^{-11}$ Hz s$^{-1}$. Such limit simply reflects the shortness of the time interval covered by the {\\xmm} observation. As a matter of fact, if a NS spinning at the frequency of {\\igr} simply gains the specific angular momentum of the accreted matter, the expected spin frequency derivative cannot be larger than \\begin{equation} \\dot{\\nu}\\simlt 2.4\\times10^{-13} I_{45}^{-1} L_{37} R_{10} m_{1,1.4}^{-1/3}\\;Hz/s. \\end{equation} Here $I_{45}$ is the NS moment of inertia in units of $10^{45}$ g cm$^2$, $m_{1,1.4}$ is the NS mass in units of 1.4 $M_{\\odot}$, $R_{10}$ is the NS radius in units of 10 km, and $L_{37}$ is the bolometric X-ray luminosity in units of $10^{37}$ erg s$^{-1}$, the dependence on which have been introduced using the relation $L_X= GM\\dot{M}/R_{NS}$.  Thus, considering the estimate of $L_X$ we give from spectral modelling ($\\simeq 1.6\\times 10^{37}$ d$_8^2$ erg s$^{-1}$) the maximum expected spin frequency derivative is nearly two order of magnitude lower than the loose upper limit we could set, given the length of the {\\xmm} observation.  The pulsed fraction is among the largest ever observed from an AMSP (14.4(3) \\% RMS). Typical observed values are between 2 and 8 \\%, with sporadic increases up to $\\sim$ 10\\% especially at the end of the outburst episodes \\citep{WjnvdK98,Mark02,Gal02,Cam03,Wts05,Gal05,Gal07,Cas08,Alt08,Pat10}, while \\citet{Pat09b} detected a pulse amplitude up to 19\\% from XTE J1807--204.  If the misalignment angle between the magnetic dipole and the spin axis is small ($\\theta\\simlt 20^{\\circ}$), the observed value of the pulsed fraction suggests that the inclination is $\\simgt 45^{\\circ}$ \\citep{PouBel06}. Nevertheless, the degree of anisotropy of the emitted light and the assumed spot shape may significantly influence the pulse amplitude. A detailed analysis of the pulse profile is mandatory to derive firm constraints on the spot geometry, and as it is beyond the scope of this paper, we defer it to another work. It is worth to note, however, that the presence of a hump in antiphase with respect to the global maximum possibly suggests that we see, at least for a fraction of rotational phases, the antipodal spot \\citep[see e.g.][]{Leh09}.  \\begin{figure} \\begin{centering} \\includegraphics[angle=270,width=\\columnwidth]{F9.eps} \\caption{Fractional contribution of the \\texttt{bbodyrad} and \\texttt{compps} components to the EPIC-pn countrate (top curve, left scale), and first harmonic fractional amplitude (bottom curve, right scale). } \\label{fig:ampl_spectrum} \\end{centering} \\end{figure}  The pulse profile of {\\igr} shows strong spectral variability, similarly to what has been observed from the AMSPs, {\\xte} and {\\saxj}, by GP05 and \\citet{Pat09}, respectively. The decrease of the pulsed fraction at low energies can be completely ascribed to to the growing influence of the accretion disc emission. To see this, we plot in Fig.\\ref{fig:ampl_spectrum} the fractional contribution of the \\texttt{bbodyrad} and the \\texttt{compps} component (Model PS-refl), to the overall EPIC-pn countrate, $f_{(BB+PS)}$ (top curve, right scale). The first harmonic fractional amplitude is also plotted (bottom curve, right scale, same as Fig.\\ref{fig:ampiezze}).  The similarity between the shape of the two curves is striking. The correlation can also be quantitatively expressed, as the ratio between the amplitude in the range 1--2 keV and the maximum amplitude ($A_1^{max}=0.232(7)$, E=9--11 keV) is 0.71(2), almost equal to $f_{(BB+PS)}$ in the same energy band (0.675(3)). The energy dependence of the pulsed fraction therefore represents a compelling evidence that the bulk of the emitted spectrum is somewhat related to the magnetic caps, where the pulsations are formed.  The similarities with the pulse amplitude energy dependence of other AMSPs (see above), further strengthens such a conclusion.  Soft phase lags of $\\sim230\\mu s$ have also been observed. Soft lags of a similar amount ($\\sim 200\\mu s$) were observed from the AMSP {\\saxj} by \\citet{cui98}, who interpreted them in terms of Compton downscattering of intrinsically hard photons in a $\\tau\\sim 10$ cloud, which is however not observed in the X-ray spectrum. It was also proposed that soft lags could be due to the Doppler energy shifts introduced by the NS fast rotation \\citep{Frd00,Wnb01}, or due to the different angular distribution of the flux emitted by the blackbody and the Comptonized emission \\citep{GrlDonBar02,PouGie03}. As the lags we observe from {\\igr} are of similar amplitude than those of {\\saxj} (which spins 1.7 times faster than this source), Doppler effects are unlikely the only physical reason behind phase lags.  \\subsection{The companion star} \\label{sec:companion}  {\\igr} is one of the AMSP with the longest orbital period ever observed. From the measured mass function, $f(m_2;m_1,i)=(m_2\\sin{i})^3/(m_1+m_2)^2$ (see Table \\ref{table:spin}), it is possible to derive constraints on the mass of the companion star, $m_2$. Here $f$,$m_1$ and $m_2$ are expressed in solar masses. The absolute minimum on $m_2$ is reached considering an inclination of 90$^{\\circ}$, which translates in $m_2^{(min)}\\simeq0.05+0.06m_1$ for $1\\leq m_1 \\leq 2$. The absence of eclipses limits the range of viable inclinations to $i<90^{\\circ}-\\gamma$, where $\\gamma=\\arctan{(R_{2}/a)}$ is the angle subtended by the companion star as seen from the NS, $R_2$ is the radius of the companion star and $a$ is the binary separation. In order for mass transfer to proceed the radius of the companion must be comparable with the radius of the Roche Lobe, $R_2\\simlt R_{L2}$, which is estimated through the relation given by \\citet{Pcz71}: \\begin{equation}R_{L2}=0.462 a [m_2/(m_1+m_2)]^{1/3}.\\label{eq:pcz}\\end{equation}  To evaluate the maximum inclination compatible with the absence of eclipses, we calculate $\\gamma$ for $R_2=R_{L2}(m_2^{(min)})$, obtaining $i\\simlt 78^{\\circ}$.  The absence of dips in the X-ray light curve further decreases the range of possible inclinations, so that we consider $i\\simlt70^{\\circ}$ only, slightly increasing the minimum companion mass to, e.g., $m_2\\simgt0.15$, for $m_1=1.4$. A main sequence companion star is clearly indicated. To see this we use the third Kepler law in Eq.(\\ref{eq:pcz}) to relate the Roche Lobe radius to the mean density of the companion star at the observed orbital period: \\begin{equation} r_{L2}(m_2)=R_{L2}/R_{\\odot}=0.536\\;  m_2^{1/3}. \\end{equation} This relation is plotted in Fig.\\ref{fig:mass2}, together with the Zero Age Main Sequence (ZAMS) mass-radius relation of \\citet{ChbBrf00}, $r_2^{CB}(m_2)$. The two curves cross for m$_2^{(max)}=0.44$ and r$_2^{(max)}=0.41$. Heavier companion stars can be safely excluded as the Roche Lobe would be overfilled, and mass transfer unstable. Assuming $m_1=1.4$, this solution corresponds to i$^{(min)}=20.5^{\\circ}$.   \\begin{figure} \\includegraphics[angle=270.0,width=\\columnwidth]{F10.eps} \\caption{Size of the companion Roche Lobe (top curve) and ZAMS mass radius relation as from \\citet[][bottom curve]{ChbBrf00}. The values taken by an assumed Roche Lobe filling companion star, for a set of inclinations and $m_1=1.4$, are also indicated. The left shaded region is excluded from the absence of eclipses, while the right shaded one from the assumption that the companion is not larger than its Roche Lobe. The minimum mass for $i< 70^{\\circ}$ is also plotted with a vertical dashed line. } \\label{fig:mass2} \\end{figure}   The reflection continuum and the iron K$\\alpha$ line we have modelled with a \\texttt{diskline} indicate an inclination in the range 38--68$^{\\circ}$ (see Table \\ref{tab2} and Sec.\\ref{disc:spectrum}). This interval corresponds to $m_2=$0.15--0.23 for m$_1=1.4$. Slightly larger values are found for heavier NS.  The inclination estimate we have drawn from spectral fitting thus indicates that the companion is slightly bloated with respect to its ZAMS thermal equilibrium radius.  A possible mechanism to drive a companion star out of thermal equilibrium is irradiation by the compact object, but to what extent it is important is still matter of debate \\citep[see][and references therein]{Rtr08}. However, irradiation is not strictly needed to account for the values we observe from {\\igr}, as CV-like evolution can explain it if the companion star is slightly evolved at the onset of mass transfer (but still before the MS turn off), or heavier than 1 M$_{\\odot}$. As \\citet{PylSvn89} shown, angular momentum losses driven by magnetic braking and gravitational radiation can still lead to a converging system (that is, a system that evolves decreasing its orbital period), even for slightly evolved, or heavy, companions. In these cases, the companion radius will be systematically larger than the ZAMS radius. To see this, we consider the evolutionary tracks calculated by \\citet{PdsRppPfh02}. While an unevolved companion with initial mass of $m_2^{(i)}$=1 crosses the observed orbital period of {\\igr} when m$_2\\simeq0.44$, as its radius is frozen to the ZAMS value, this happens for $m_2\\simeq0.16$ if it is slightly evolved, and thus larger, at the beginning of mass transfer. Similar results are obtained even for initially unevolved, but heavier, companion star (the evolutionary track for a $m_2^{(i)}=1.4$ companion star passes through the period of {\\igr} when $m_2=0.15$). Values of $m_2$ in agreement with the inclination indicated by spectral modelling can be thus obtained, for peculiar conditions of the companion star at the onset of mass transfer.  \\subsection{The X-ray spectrum} \\label{disc:spectrum}  We presented in Sec.\\ref{sec:xmmspectrum} and \\ref{sec:rxtexmm} a detailed spectral analysis of the 0.5--200 keV X-ray emission of the AMSP {\\igr}, using two different Comptonization models, and checking the presence reflection.  The spectrum of {\\igr} is modelled by four components, (i) the accretion disc emission peaking at $\\simeq 0.3$ keV; (ii) a $\\simeq$ 0.6 keV blackbody of apparent radius $\\simeq7$ d$_8$ km; (iii) Comptonization from a hot ($kT_e\\simeq50$ keV) medium of moderate optical depth ($\\tau\\simeq1.3$), of the thermal ($\\simeq1.3$ keV) photons provided by a $\\simeq3.5$ d$_8$ km region (iv) Compton reflection and an iron K$\\alpha$ emission line, interpreted as coming from the accretion disc illuminated by the hard radiation emitted from the NS surface.  The shape of the best fitting model does not change much using different Comptonization models.  As the model PS-refl better addresses a number of physical aspects of this source emission (see Sec.\\ref{sec:rxtexmm}) we discuss the parameters thus obtained. Compton disc reflection is significantly detected regardless of the particular model used to describe Comptonization, as well as an iron line at an energy compatible with Fe XXV-XXVI.  Even if reflection has not always been observed, such a spectral decomposition have already proved successful in the description of the X-ray spectrum AMSP (\\citealt{GrlDonBar02}, GP05, \\citealt{Fal05,Pap09,Pat09}). In particular, the presence of disc thermal emission has always been detected when {\\xmm} high resolution spectra extending down to $\\sim$1 keV were available.  The similarities in the X-ray emission of AMSP can be viewed as an indication of how the physical processes that produce the bulk of their spectra are similar, and probably related to the presence of a magnetosphere in these systems.  \\label{disc:column}  The correlation between the energy dependence of the pulsed fraction and the spectral decomposition we have employed strongly suggests that the Comptonizing medium surrounds the hot spots on the NS surface. \\citet{GrlDonBar02}, \\citet{PouGie03}, GP05, \\citet{Fal05} and \\citet{Fal05b} interpreted the similar hard components shown by other AMSPs as coming from Comptonization in a plasma heated in the accretion columns. The magnetic field collimates the in-falling matter in columns of radii that are few tenths of the NS size. Even for accretion rates $\\sim 0.1$ $\\dot{M}_{Edd}$ like those observed for these sources during outbursts, the local accretion rate $\\dot{m}$ can therefore attain the Eddington level \\citep{BasSun76}.  In such a case, the radiation pressure becomes comparable to the ram pressure of the plasma falling at supersonic velocities, and a shock may thus form, heating the electrons up to the observed large temperatures.   The observed $\\simeq 0.6$ keV blackbody component is too cold to supply enough seed photons for the observed Comptonized spectrum. Disentangling the temperature of seed photons from that of the observed blackbody significantly improves the fit and ensure the energy balance between the hot plasma and the region that provides the seed photons. A similar result has already been obtained by GP05, modelling the X-ray spectrum of the AMSP, {\\xte}.  The size of the region that provides the soft photons, $R_{soft}^{\\infty}=3.4^{+0.1}_{-0.4}$ d$_8$ km, is compatible with the expected radius of an hot spot.  The radius of the observed blackbody, $R_{BB}^{\\infty}=6.3^{+0.3}_{-0.1}$ d$_8$ km, is instead of the same order of the asymptotic value taken by the radius of the thermal burst emission. This component is therefore emitted by a larger fraction of the NS surface.  \\label{sec:disc} The cooler thermal component can be safely attributed to the emission of an optically thick accretion disc. The contribution of this component at low energy is responsible for the decrease of the pulsed fraction. Moreover, the parameters we obtain with a \\texttt{diskbb} model nicely fit the expectations for an accretion disc around an AMSP. Considering the range of inclinations indicated by the reflection component (38--68$^{\\circ}$), the model \\texttt{diskbb} evaluates an apparent inner disc radius $R_{in}$ in the range 18--28 d$_8$ km. Such a radius is estimated approximating the disc temperature profile with its asymptotic behaviour $T(r)\\propto r^{-3/4}$, without accounting for the decrease of the viscous torque at the inner boundary of the disc, and neglecting spectral hardening. To evaluate the importance of these effects we consider the disc model \\texttt{diskpn} \\citep{Grl99}. The spectral shape found by \\texttt{diskbb} is recovered after rescaling the inner disc radius $R_{in}$ by a factor $\\approx2.3$, for $m_1=1.4$, $d=8$kpc, $i=50^{\\circ}$ and $f=1.7$, where f is the ratio between the colour and effective temperature \\citep{ShmTkh95}.  The inner radius indicated by disc emission modelling thus meets the request that an accretion disc around an X-ray pulsar is not truncated at radii much larger than the corotation radius, $R_{C}=(GM/4\\pi^2\\nu^2)^{1/3}=42.8 \\: m_{1.4}$ km for {\\igr}, where $m_{1.4}$ is the NS mass in units of 1.4 $M_{\\odot}$. Also the ratio between the flux in the disc component and the total flux agrees with a truncation radius a few tens of km away from the NS star. From the virial theorem it is in fact expected that $F_{disc}/F_{tot}\\simeq R_*/2r_{in}$. As we measure $F_{disc}/F_{tot}\\simeq0.1$, the flux ratio predicts $r_{in}\\simeq5R_*$, which substantially agrees with the measured value.  The presence of reflection features, such as an iron K$\\alpha$ emission line and Compton reflection, is significantly detected in spectral fitting.  The iron line stands at $\\sim 4\\sigma$ above the continuum and has an equivalent width of $\\simeq40$ eV. If modelled with a Gaussian, its width can only be constrained to be $\\simgt 0.5$ keV. As the spectrum of this source is dominated by Comptonization in a hot and optically thin medium, it is hard to find alternatives to disc reflection in order to explain such a width. In the inner regions of the accretion disc the line is broadened and red-shifted by the effects of Keplerian motion and by the gravitational influence of the nearby NS. Modelling the feature with a \\texttt{diskline} we obtain an estimate of the inner disc radius as $r_{in}=27^{+6}_{-9}$ R$_g$, which translates in $r_{in}=56^{+12}_{-19}$ km for a 1.4 M$_{\\odot}$ NS. Even if it is rather loose due to the limited counting statistics, also this estimate overlaps with the corotation radius, as it is expected from accretion theories. The addition of the {\\rxte} dataset to the {\\xmm} spectrum also allows the detection of the hump at $\\sim 30$ keV expected from Compton reflection.  The ionised state of the reflector is rather high, $\\log\\xi=3.0^{+0.4}_{-0.2}$.  It is worth to note that the highly ionised surface in the reflector is in agreement with the transition energy of the broad iron line at E$_{Fe}=6.82^{+0.09}_{-0.11}$ keV, that we identify as a Ly$\\alpha$ resonant transition of Fe XXV (rest-frame energy at 6.70 keV), that is likely to be produced in a photo-ionised plasma at $\\log\\xi\\simeq 3$ \\citep{Klm04}.  Considering the values indicated by the reflection continuum and by the iron line modelling, we give a conservative estimate of the inclination in the range 38--68$^{\\circ}$.  Also the OVIII edge we find at a high absorption depth ($\\tau=$0.1--0.3) can be tentatively interpreted in terms of disc reflection. As its width is $<34$ eV, the outer rings of the disc should be involved, in order to make the rotational broadening negligible.  \\subsection{The X-ray bursts} \\label{disc:burst}  In Sec. \\ref{sec:burst} we have presented a spectral and temporal analysis of the bursts observed by {\\xmm}. Considering the {\\it Swift} observations, \\citet{Bozzo09} concluded that the observed recurrence times could be reconciled with the {\\it persistent} flux if it is assumed that the bursts are ignited in a pure helium environment.  This hypothesis can be checked estimating the local accretion rate $\\dot{m}$ needed to produce the fluence of the second burst, which lags the first by $\\Delta t=$40.9 ks.  One has in fact that $\\dot{m}=y/\\Delta t$, where $y$ is the column depth at which the burst is ignited. The value of $y$ can be estimated as $y=4\\pi d^2 \\mathcal{F}/Q_{nuc} R_*^2$ \\citep[see e.g.][]{Glw08}, where $\\mathcal{F}=(1+z)\\tau F_{peak}^{\\infty}$ is the burst fluence, and $Q_{nuc}$ is the energy per nucleon released during the thermonuclear burning ($Q_{nuc}\\approx 1.6$ MeV for complete burning of He into iron group elements,\\citealt{WalWoo81}). The measured values (see Table \\ref{table:burst}) yield $y=(2.6\\pm0.6)\\times10^8$ d$_8^2$ R$_{10}^{-2}$ g cm$^{-2}$, and $\\dot{m}=(6.3\\pm1.4)\\times 10^3$ d$_8^2$ R$_{10}^{-2}$ g cm$^{-2}$ s$^{-1}\\simeq 0.08\\;R_{10}^{-1} \\dot{m}_{Edd}$, where $z=1.31$ is considered. This value of $\\dot{m}$ agrees with the bolometric luminosity of the {\\it persistent} emission. Assuming that the \\texttt{blackbody} and the Comptonized component arise from the vicinity of the NS surface, and therefore the observed fluxes have to be corrected for general relativity effects, we in fact estimate from model PS-refl the {\\it persistent} luminosity as, $L_X^{pers}=4\\pi d^2[(1+z)^2(F_{BB}+F_{PS})+(F_{Disc}+F_{Refl})]=(1.6\\pm0.1)\\times 10^{37}$ d$_8^2$ erg s$^{-1}\\;\\simeq0.09$ m$_{1.4}^{-1}$ L$_{Edd}$. The agreement between the {\\it persistent} luminosity and the mass accretion rate needed to explain the bursts fluence and recurrence time, when it is assumed that all the hydrogen is steadily burnt between bursts, suggests that bursts are triggered in nearly pure helium environment.  It is perhaps useful to clarify that the difference between the local accretion rate estimated by the burst properties, $\\dot{m}\\simeq 0.08 \\dot{m}_{Edd}$, and the rate needed for a shock to form in the accretion column, $\\dot{m}\\simeq\\dot{m}_{Edd}$, arises as the first is evaluated averaging the accretion rate over the whole NS surface, while the latter considers only the column size. Such a treatment seems appropriate as a magnetic field $\\simlt 10^9$ G like the one expected in an AMSP should be able to confine accreted matter in the polar caps up to a column density of $\\simeq 10^7$ g cm$^{-2}$ \\citep{BrwBld98}, which is well below the value we measure. The accreted layer is therefore able to diffuse across the whole NS surface before the burst onset.   As no photospheric radius expansion is observed, an upper limit on the distance can be set as $d^2\\simlt{L_{PRE}}/{4\\pi F_{peak}^{\\infty}}$, where the luminosity at which photospheric radius expansion is expected, $L_{PRE}=(3.79\\pm0.15)\\times10^{38}$ erg s$^{-1}$, was empirically derived by \\citet{Kulk03}. Considering the peak flux of the second burst and the associated uncertainty, we set an upper limit on the distance to the source of $10$ kpc at a 90\\% confidence level. The radius of the thermal emission attained during the burst decay can be identified with the NS radius, and used to put a reasonable lower limit on the distance to the source. The measured value $R_{app}=(7.0\\pm0.3) d_8$ km can be related to the effective radius through the relation, $R=R_{app} f_C^2(1+z)^{-1}$ \\citep[see e.g.][]{LewvPrTam93}. The factor $f_C$ takes into account spectral hardening, and has been estimated as $f_C\\simeq1.35$ by \\citet{Mdj04}, for the case of a source that does not reach the Eddington level. Solving the previous relation for $M>1.4$ M$_{\\odot}$, and considering that a soft equation of state like the A predicts $R_*>8.5$ km and M$\\simlt 1.92$ M$_{\\odot}$ for a NS spinning at the rate of {\\igr} \\citep{Pan71,CokShpTeu94}, we obtain $d\\simgt 6.5$ kpc, unless a strange star is considered. For a lighter NS (M=1.57 M$_{\\odot}$) this limit increases to 6.9kpc. Given the estimated range of distances, it is possible to safely state that {\\igr} belongs to the Galactic bulge. Also the compatibility of the measured equivalent hydrogen absorption column with that observed from {\\xte} (whose distance was estimated to be $\\simgt$ 7 kpc, \\citealt{Mark02,Pap08}) supports such a conclusion.   Pulsations are detected throughout the burst decay at a similar amplitude than that seen during the {\\it persistent} emission, while during the first few seconds they disappear, probably because of the telemetry limitations of the EPIC-pn.  Their phase is also compatible within 0.1--0.2 cycles with that of pre-burst pulsations, suggesting that magnetic channelling of accreted matter and the burst onset happen at a similar location on the NS surface. This indication is strengthened by the analysis performed by Riggio et al. (2010, submitted) over RXTE data. They find in fact that, with the exception of the first two seconds since the burst onset, the burst oscillations are phase locked to pre-burst pulsations within 0.1 cycle.  Phase locking between burst and non-burst oscillations of an AMSP has already been observed by \\citet{Wts08}, at a much larger accuracy ($\\simeq$0.01 cycles) than that made available by the data presented here.  As they discuss (see also references therein), the most appealing interpretation is in terms of the effect a temperature gradient between the fuel impact point on the surface, and the rest of the surface, may have on the ignition conditions, or on its development. A temperature gradient during the {\\it persistent} emission is indeed indicated by our spectral modelling, which requires the region feeding the accretion columns of soft photons to be hotter than the rest of the surface. A firm explanation of the correlation between pulsations observed during bursts and during the {\\it persistent} emission is anyway still missing."
    },
    "1005/1005.1633_arXiv.txt": {
        "abstract": "\\noindent Using the new Cosmic Origins Spectrograph (COS) on the {\\it Hubble Space Telescope (HST)}, we obtained moderate-resolution, high signal/noise ultraviolet spectra of HD 209458 and its exoplanet HD 209458b during transit, both orbital quadratures, and secondary eclipse. We compare transit spectra with spectra obtained at non-transit phases to identify spectral features due to the exoplanet's expanding atmosphere. We find that the mean flux decreased by $7.8\\pm 1.3$\\% for the C~II 1334.5323~\\AA\\ and 1335.6854~\\AA\\ lines and by $8.2\\pm 1.4$\\% for the Si~III 1206.500~\\AA\\ line during transit compared to non-transit times in the velocity interval --50 to +50 km~s$^{-1}$. Comparison of the C~II and Si~III line depths and transit/non-transit line ratios shows deeper absorption features near --10 and +15 km~s$^{-1}$ and less certain features near --40 and +30--70 km~s$^{-1}$, but future observations are needed to verify this first detection of velocity structure in the expanding atmosphere of an exoplanet. Our results for the C~II lines and the non-detection of Si~IV 1394.76~\\AA\\ absorption are in agreement with \\citet{Vidal-Madjar2004}, but we find absorption during transit in the Si~III line contrary to the earlier result. The $8\\pm 1$\\% obscuration of the star during transit is far larger than the 1.5\\% obscuration by the exoplanet's disk. Absorption during transit at velocities between --50 and +50~km~s$^{-1}$ in the C~II and Si~III lines requires high-velocity ion absorbers. Assuming hydrodynamic model values for the gas temperature and outflow velocity at the limb of the outflow as seen in the C~II lines, we find mass-loss rates in the range (8--40)$\\times 10^{10}$ g~s$^{-1}$. These rates assume that the carbon abundance is solar, which is not the case for the giant planets in the solar system. Our mass-loss rate estimate is consistent with theoretical hydrodynamic models that include metals in the outflowing gas. ",
        "introduction": "\\citet{Mazeh2000} monitored the radial velocities of the G0~V star HD~209458 to derive the orbital parameters and minimum mass of its transiting planet HD 209458b. Very shortly thereafter, accurate measurements of the mass and radius of HD~209458b were extracted from radial velocity and photometric transit observations by \\citet{Henry2000} and \\citet{Charbonneau2000}. Table 1 lists the presently accepted properties of HD~209458 and HD~209458b cited by \\citet{Knutson2007}. HD~209458b is likely the best-studied Jupiter-like exoplanet located very close to its host star. As a result of its proximity, HD~209458b receives very strong incident radiation (see estimate in Table~1) and stellar wind flux from its host star.  HD 209458b is the first transiting planet for which atmospheric absorption was observed in the resonance lines of Na~I \\citep{Charbonneau2002} and subsequently in lines of H~I, C~II, O~I, and others. Using the G140M grating of Space Telescope Imaging Spectrograph (STIS) on the Hubble Space Telescope {\\it HST} with its resolution of $\\sim 30$ km~s$^{-1}$, \\citet{Vidal-Madjar2003} found that the Lyman-$\\alpha$ line flux was reduced by $15\\pm 4$\\% during transit at velocities between --130 and +100 km~s$^{-1}$. Based on this first detection of H~I absorption at or above the planet's Roche lobe, they concluded that hydrogen is escaping from the planet. \\citet{Vidal-Madjar2004} then used the low-resolution G140L grating of STIS to detect an absorption depth at midtransit of $5\\pm 2$\\% for the unresolved Lyman-$\\alpha$ line, $13\\pm 4.5$\\% for the O~I 1304~\\AA\\ multiplet, and $7.5\\pm 3.5$\\% for the C~II 1335~\\AA\\ doublet. Their detection of absorption in the O~I and C~II lines at or above the Roche lobe confirmed significant mass loss from the planet by hydrodynamic outflow. They argued that reduction in the Lyman-$\\alpha$ line flux at $\\approx 100$ km~s$^{-1}$ from line center indicates high velocities of neutral hydrogen atoms. Additionally, more absorption on the blue side of the Lyman-$\\alpha$ emission line than the red side indicates outflow towards the observer and away from the planet. A subsequent observation of HD~209458 using the {\\it HST} Advanced Camera for Surveys (ACS) showed a reduction in the unresolved Lyman-$\\alpha$ line flux of ($8.0\\pm 5.7$)\\% during transits, consistent with the STIS G140L and G140M results, thereby confirming mass loss from the planet's large exosphere \\citep{Ehrenreich2008}.  There have been a number of reanalyses of the data and futher discussions of the STIS transit G140M observations by \\citet{Ben-Jaffel2007}, \\citet{Ben-Jaffel2008}, \\citet{Vidal-Madjar2008} and others. There is a consensus that HD~209458b has an extended exosphere and is losing mass, but important details are not yet understood. Such questions as: whether the exosphere size is larger or smaller than the projected Roche lobe, what are the outflow speed and mass-loss rate, and whether or not the mass loss from the planet produces a cometlike tail are not yet answered with the existing STIS and ACS observations. While the STIS G140M Lyman-$\\alpha$ spectrum obtained by \\citet{Vidal-Madjar2003} provides some velocity information, higher spectral resolution and signal/noise (S/N) are needed to obtain the velocity structure and optical depth of the outflowing gas in Lyman-$\\alpha$ and other spectral lines.  The large abundance of hydrogen and large flux in the Lyman-$\\alpha$ line make this line a prime candidate for studying exoplanet outflows, but there are uncertainties (many of them unique) in measuring the exoplanet's mass-loss rate only from this line. These include the modest S/N of the existing STIS G140M grating data, possible variability of the stellar Lyman-$\\alpha$ profile at the time of transit that cannot be directly measured, broad interstellar absorption in the core of the line that prevents measurements of low velocity absorption during transit, and the time-varying geocoronal and interplanetary emission in the line core that cannot be completely corrected for.  Absorption of Lyman-$\\alpha$ photons during transit over the velocity range --130 to +100 km~s$^{-1}$ raises the question of the origin of the high-velocity hydrogen that is difficult to explain by a thermal plasma at moderate temperature. \\citet{Holmstrom2008} proposed that the high-velocity neutral hydrogen can be produced  by charge exchange between stellar wind protons and neutral hydrogen in the planet's outflow. Such energetic neutral atoms (ENAs) are seen in the solar system and likely occur where the stellar and planetary winds interact between HD~209458b and its host star. If this were the only possible explanation, then the Lyman-$\\alpha$ profile during transit provides useful information on the stellar wind but ambiguous information on mass loss from the planet \\citep{Holmstrom2008, Murray-Clay2009}. However, \\citet{Lecavelier2008b} argued that the observed Lyman-$\\alpha$ profile and the high-velocity hydrogen atoms can be simply explained by stellar radiation pressure and the ENA explanation required a peculiar stellar wind model.  The Lyman-$\\alpha$ transit observations by \\citet{Vidal-Madjar2003, Vidal-Madjar2004} stimulated a number of interpretive papers concerning the mechanisms for mass loss and lifetime of a Jupiter-like planet close to its host star [e.g., \\citet{Vidal-Madjar2003, Baraffe2004, Koskinen2007, Holmstrom2008, Koskinen2010}] and hydrodynamic models [e.g., \\citet{Lammer2003, Lecavelier2004, Yelle2004, Jaritz2005, Tian2005, Garcia2007, Schneiter2007, Murray-Clay2009}]. As discussed below, the theoretical models differ in their treatment of the amount of available radiative energy input (photoionization and otherwise) to drive the hyrdodynamic blow-off, the location of the heating, the presence of metals in the outflow, and the relative amount of cooling by expansion, thermal conduction, and radiation. These models predict very different mass-loss rates. Which of these models are consistent with the spectrally resolved metal lines observed during transit?  With the objectives of measuring an accurate mass-loss rate from the planet and testing different physical models for the mass loss, one needs far better data, which requires a new scientific instrument. One needs higher S/N spectra with higher spectral resolution and sufficient sensitivity to study many spectral lines formed in the extended atmosphere and wind of the exoplanet. In particular, it is important to use spectral lines other than Lyman-$\\alpha$ to avoid unique difficulties in the analysis of this line. As we describe below, the new Cosmic Origins Spectrograph (COS) on {\\it HST} is well designed for this task. ",
        "conclusions": "The high sensitivity and moderate spectral resolution of COS allowed us to obtain the first measurements of gas absorption in the exosphere of a transiting exoplanet in lines of C~II and Si~III with velocity resolution. During transits of HD~209458, we found that the mean flux in these lines is reduced by $7.8\\pm 1.3$\\% for the co-added C~II lines and $8.2\\pm 1.4$\\% for the Si~III line in the velocity range --50 to +50~km~s$^{-1}$. The 8\\% absorption constraint sets the location of the outflow limb (2.36R$_{planet}$) where the C~II lines become optically thin. Theoretical hydrodynamic models show that the outflow velocity is about 10 km~s$^{-1}$ and gas temperature about 10,000~K at this location. If we assume that the outflow is purely thermal, then $\\dot{M}_{total}=8\\times 10^{10}$ g~s$^{-1}$, somewhat larger than theoretical estimates based on outflows with no metals. Since absorption in the C~II lines is much broader than indicated by Doppler broadening with a thermal velocity of 3.7 km~s$^{-1}$, the lines must be broadened by large turbulence or high-velocity C$^+$ ions. If turbulent broadening dominates, then the 8\\% absorption seen out to $\\pm 50$ km~s$^{-1}$ would require mass-loss rates as large as $1\\times 10^{12}$ g~s$^{-1}$, but such highly supersonic velocities are not likely. Instead, absorption by high-velocity C$^+$ and Si$^{++}$ ions formed by the interaction of the stellar wind and planetary mass loss or some other mechanism broadens the C~II and Si~III lines. We conclude that the mass-loss rate in the exoplanet's outflow is in the range (8--40)$\\times 10^{10}$ g~s$^{-1}$. These rates assume that the carbon abundance is solar, which is not the case for the giant planets in the solar system. We note that mass-loss rates predicted by hydrodynamic models that include solar abundant metals in the outflow lie at the upper limit of this range.  Comparison of the C~II and Si~III line depths and transit/non-transit line ratios shows deep absorption features near --10 and +15 km~s$^{-1}$ in both lines. Since the C~II and Si~III lines are likely formed in the same region of the exosphere, we include in Figure~4 the co-addition of the two profiles with typical error bars for comparison with theoretical models. The increase in S/N with the co-addition of the C~II and Si III spectra allows us to make more definitive statements about the velocity structure. The dips near --10 and +15 km~s$^{-1}$ are more than twice as large as the errors per spectral resolution element (about 17 km~s$^{-1}$) and are most likely real. In addition, the ratios near --40 and +30--70 km~s$^{-1}$ are also low, but the line fluxes decrease rapidly away from line center and these low ratios are less certain. Future observations with higher S/N are needed to confirm this first detection of velocity structure in the expanding atmosphere of an exoplanet.  The COS transit observations of the absorption velocity structure in these lines demonstrate the need for including high-velocity C$^+$ and Si$^{++}$ ions in models of the exosphere of HD~209458b and presumably other close-in exoplanets. It remains to be seen whether the inclusion of a simple distribution of ions with high effective temperature as proposed by \\citet{Ben-Jaffel2010} is adequate to explain the data. The new COS data provide the opportunity to test a range of  absorption models based on such physical mechanisms as radiation pressure, charge exchange, and interactions between the stellar wind and the planet's outflow to understand mass-loss processes in the exospheres of close-in planets.   This work is supported by NASA through grants NNX08AC146 and NAS5-98043 to the University of Colorado at Boulder. We thank the many people at the Space Telescope Science Institute, Goddard Space Flight Center, Ball Aerospace and Technologies Corp. and CASA at the University of Colorado for the excellent hardware and software that has made COS a great success. We particularly wish to thank the referee for his many thoughtful and thorough comments and his bringing to our attention important publications that have led to significant modifications in the paper. We thank Dr. J. Schneider for encouraging us to look for an assymetry between ingress and egress. We also thank Dr. Lofti Ben-Jaffel for clarifying our understanding of the Ben-Jaffel \\& Hosseini (2010) paper. We thank Dr. Alain Lecavelier for calling our attention to the different values of $\\tau$ at the limb that depend on the density structure in the envelope."
    },
    "1005/1005.0635_arXiv.txt": {
        "abstract": "We present {\\it Chandra X-ray Observatory} archival observations of the supernova remnant 1E0102.2-7219, a young Oxygen-rich remnant in the Small Magellanic Cloud.  Combining 28 ObsIDs for 324\\,ks of total exposure time, we present an ACIS image with an unprecedented signal-to-noise ratio (mean S/N $\\simeq$ $\\sqrt{S}$ $\\sim$ 6; maximum S/N $>$ 35) .  We search within the remnant, using the source detection software {\\sc wavdetect}, for point sources which may indicate a compact object.  Despite finding numerous detections of high significance in both broad and narrow band images of the remnant, we are unable to satisfactorily distinguish whether these detections correspond to emission from a compact object.  We also present upper limits to the luminosity of an obscured compact stellar object which were derived from an analysis of spectra extracted from the high signal-to-noise image. We are able to further constrain the characteristics of a potential neutron star for this remnant with the results of the analysis presented here, though we cannot confirm the existence of such an object for this remnant. ",
        "introduction": "The supernova remnant 1E 0102.2-7219 (SNR 1E0102.2-7219, hereafter, ``E0102'') was first identified by \\citet{SM81} in a survey of the Small Magellanic Cloud (SMC) with the {\\it Einstein X-Ray Observatory}\\@. While this oxygen-rich remnant resides at a distance of 59\\,kpc \\newline\\citep{VDB99}, it is far less extincted than many Galactic remnants due to its high galactic latitude ({\\it l} $\\sim$ --45$^{\\circ}$).  Kinematic studies in the X-ray and optical have determined E0102 to be a relatively young remnant, with an expansion age between 1000 and 2000 years \\citep[][respectively]{H00,Fink06}\\@. For these reasons, E0102 has been extensively studied across the electromagnetic spectrum $\\!\\colon\\!$ radio \\citep{AB93}, IR \\citep{S05}, optical \\citep{TD83, B00, Fink06}, UV \\citep{B89}, and X-ray \\citep{H94, G00, R01, F04}\\@.  E0102 was likely generated by a Type Ib or Ic core-collapse supernova event \\citep{B00}\\@.  As such, the remnant E0102 would be expected to host some variety of compact object (i.e., stellar mass black hole or degenerate stellar object). Furthermore, \\citet{AB93} and \\citet{Fink06} have considered the possibility that such a compact object may be influencing the complex morphology of this remnant.  Though this remnant has been well studied, few groups have provided observational constraints to the presence of a compact object for E0102.  In fact, only \\citet{G00} have presented an upper flux limit on a Crab-like pulsar for this remnant.  In this research, we seek to constrain the X-ray luminosity of this unidentified compact object.  E0102 is a calibration target for the {\\it Chandra X-Ray Observatory} facility team, and has been observed since {\\it Chandra's} 1999 launch.  As a result, the {\\it Chandra} Data Archive contains 2.16 million seconds of calibration observations acquired during 215 observing sessions (ObsIDs).  Many of these observations were obtained to calibrate the off-axis performance of the detector.  We use archive calibration data exclusively in this research, specifically those ObsIDs with high spatial resolution acquired on-axis between 2000 and 2009.  With these combined data, we present new upper limits to the X-ray flux of a compact stellar object for the remnant. ",
        "conclusions": "We have combined 324\\,ks of {\\it Chandra} X-ray observations of the SNR 1E0102.2---7219 in the SMC to produce the highest signal-to-noise image of this young, oxygen-rich remnant.  In this research, we attempt to identify a compact stellar object for this remnant within a generous search area motivated by observations of this remnant, Galactic SNRs, and compact stellar objects.  We cannot confirm the presence of a compact object for this remnant, but are able to provide the strongest upper limits on the flux, assuming that the spectral behavior of a faint neutron star for this remnant is best described by a blackbody or power-law spectrum energy distribution."
    },
    "1005/1005.5010_arXiv.txt": {
        "abstract": "We describe several projects addressing the growth of galaxies and massive black holes, for which adaptive optics is mandatory to reach high spatial resolution but is also a challenge due to the lack of guide stars and long integrations. In each case kinematics of the stars and gas, derived from integral field spectroscopy, plays a key role. We explain why deconvolution is not an option, and that instead the PSF is used to convolve a physical model to the required resolution. We discuss the level of detail with which the PSF needs to be known, and the ways available to derive it. We explain how signal-to-noise can limit the resolution achievable and show there are many science cases that require high, but not necessarily diffraction limited, resolution. Finally, we consider what requirements astrometry and photometry place on adaptive optics performance and design. ",
        "introduction": "\\label{sec:intro}  The existence of tight relations between supermassive black holes (BH) and the stellar velocity dispersion, luminosity, mass, and light concentration of their host galaxies has been a major discovery of the last decade\\cite{kor95,mag98,fer00,geb00,mar03,har04,gra07}. These correlations, which imply that BHs and their host galaxies co-evolve, have resulted in a widespread observational and theoretical effort to understand their physical origin. In addition to the flow of gas inward from the host galaxy to the central regions where it can fuel a BH, a plausible mechanism capable of linking the innermost nuclear region to the whole galaxy is the feedback outward from this accreting BH (an AGN) to the host galaxy. Understanding the co-evolution of BHs and their host galaxies has become one of the most important challenges of modern astrophysics. In this section, we briefly look at 3 science topics that address this issue from different perspectives. For all of these cases, kinematics, derived from integral field spectroscopy at high spatial resolution is a key element of the analysis.  \\subsection{Black Hole Masses in nearby Active Galaxies} \\label{sec:bhmass}  \\begin{figure}[t] \\centering \\includegraphics[width=14cm]{Overlay1.eps} \\caption{Left: stellar continuum distribution and stellar kinematics in the central arcsec (80\\,pc) of NGC\\,3227. Right: increase in $\\chi ^2$ from its minimum value for Schwarzschild models as a function of black hole mass. The results for 2 inclinations (60$^\\circ$ and 75$^\\circ$) are shown. The dashed line shows the impact of including rescaled Gauss Hermite terms $h_3$ and $h_4$. Adapted from Ref~\\citenum{dav06}; see the electronic version for colour figures. } \\label{fig:science1} \\end{figure}  Measuring BH masses is a fundamental step in deriving the relations above. One of the best methods is to model the stellar kinematics, which can only be done directly for nearby galaxies. This is complementary to other methods such as reverberation mapping which is based on temporal variability of the broad emission lines, and can not only verify such masses but in principle yields information about the geometry of the broad line region. NGC\\,3227 was the first Seyfert~1 in which Schwarzschild orbit superposition modelling was applied to the stellar kinematics\\cite{dav06}. The result (see Fig.~\\ref{fig:science1}) was a 1$\\sigma$ range for the BH mass of (0.7--2.0)$\\times$10$^7$\\,M$_\\odot$. This is a robust result, although the range is rather large because of the limited signal-to-noise and also from degeneracies between the mass of the various components in the model. It is a little lower than, although just consistent with, masses derived from other methods based on reverberation mapping and X-ray variability; and suggests that NGC\\,3227 may lie a little below the $M_{\\rm BH}-\\sigma^*$ relation. This is intriguing because NGC\\,3227 is a barred galaxy, and there is some evidence that barred galaxies as a class may preferentially lie below the relation, perhaps hinting that the relations between black holes and their host galaxies are more complex than first thought.  For this work, adaptive optics is essential to reach the highest resolution possible in order to resolve the sphere of influence of the black hole; integral field spectroscopy is required to map the stellar kinematics in 2D, since the method can be unreliable with the limited information available from longslit spectroscopy; and a good estimate of the PSF is required for the modelling to ensure that beam smearing (see Sec.~\\ref{sec:kin}) is treated correctly.   \\subsection{Star Formation in Major Mergers} \\label{sec:merger}  Merging is one of the fundamental ways in which galaxies evolve and grow. Indeed, ULIRGs, the most luminous objects in the local universe, appear to be predominantly mergers whose prodigious infrared emission arises from dust heated by obscured star formation activity. At high redshift, mergers were common events and may have dominated the cosmic star formation density at $z\\sim2$, the epoch of stellar mass assembly. At the present day, however, mergers are rare. There are therefore only a few nearby examples -- of which NGC\\,6240 is one -- from which we can learn how the merging process proceeds in detail. Building on the existing qualitative picture of this merger\\cite{tec00}, we have recently used integral field spectroscopy (see Fig.~\\ref{fig:science2}) and hydrodynamical simulations to perform a quantitative analysis of the mass, luminosity, age, and origin of its young and old stellar populations\\cite{eng10}. The mass-to-light ratios indicate the bulk of the stellar mass is older than $\\sim$1\\,Gyr. Coupled with other properties (size, dispersion, etc) of the nuclei, this strongly suggests they are the progenitor bulges. Only 20--30\\% of the bolometric luminosity, and an even smaller fraction of the mass, originates in the recent starburst. The star formation history is an important issue, because it has a significant impact on the starburst scale derived from the observations. We have found that the star formation rate associated with the recent starburst has been gradually increasing over the last $\\sim$100\\,Myr -- and will continue to do so for another 100--300\\,Myr, at which point, during the final coalescence, the merger will become a ULIRG.  \\begin{figure}[t] \\centering \\includegraphics[width=12cm]{Overlay2.eps} \\caption{Stellar continuum distribution and stellar kinematics for the prototypical merger NGC\\,6240. The black hole locations\\cite{max07} are indicated. Adapted from Ref~\\citenum{eng10}; see the electronic version for colour figures. } \\label{fig:science2} \\end{figure}  The distance of this prototypical merger is such that $1\\arcsec = 500$\\,pc, and so adaptive optics is mandatory in order to spatially resolve it. Integral field spectroscopy is needed to provide a complete view of the different geometries and kinematics for the molecular gas, ionised gas, and stars. In this case, to cover the whole system requires a field of 3\\arcsec, which SINFONI can only do with its intermediate 0.05\\arcsec\\ pixel scale. This means that our resolution is pixel limited -- and so the role of adaptive optics was to provide enhanced rather than diffraction limited resolution. Independent of the resolution actually achieved, a good estimate of the PSF is required, as for the previous example, in order to treat beam smearing correctly when deriving the kinematic model.   \\subsection{Galaxy Evolution at High Redshift} \\label{sec:highz}  The highest star formation rates and greatest incidence of QSO activity occured at $z=2$--3, indicating that this was the time when galaxies were evolving most rapidly. Much effort has been put into studying galaxies at this epoch, and most of what is known about them has come from work on the `deep fields' which have been extensively surveyed at all accessible wavelengths. A recent addition to the toolbox are detailed integral field observations, making use of adaptive optics where possible\\cite{for09}. One of the key results from this work is that, although high-z galaxies are clumpy, these clumps are often part of a massive gas-rich star-forming disk -- i.e. they are not indicative of mergers. This has led to a new perpsective in which, complementary to mergers, smooth accretion of cold gas is now understood to play a major role in galaxy evolution at early cosmic time. The properties of the disks -- their sizes, rotation velocities, and masses -- have enabled scaling relations, such as the Tully-Fisher relation, to be traced out to $z\\sim2$ and have shown that these relations do evolve\\cite{cre09}, as can be seen in Fig.~\\ref{fig:science3}. Measuring these properties requires an automated fitting procedure to match a dynamical model to the 2D observations, taking into account the PSF.  Adaptive optics is an important aspect of this work, since high-z galaxies are at most 1--2\\arcsec\\ across. However, without extremely long integrations, diffraction limited observations are impractical because of the faint surface brightness of the galaxies. Thus, on 8-m class telescopes, a major role of AO is to provide seeing enhancement, to achieve resolutions of $\\sim$0.15\\arcsec\\ corresponding to about 1\\,kpc. A more crucial issue is the sky coverage: the fields from which the galaxies are selected are designed to be free of bright stars. There are very few targets with bright nearby stars, and even if a laser guide star (LGS) is used, tip-tilt correction is difficult. Fortunately, the relaxed requirements on resolution lead directly to relaxed requirements on tip-tilt; and in some cases, it has been possible to observe without tip-tilt at all.  \\begin{figure}[t] \\centering \\includegraphics[width=14cm]{Overlay3.eps} \\caption{Left: H$\\alpha$ distribution of BX\\,482 together with the observed and modelled kinematics. Right: Tully-Fisher Relation for disk galaxies at $z\\sim2$, which are offset from the local $z=0$ relation, indicating evolution in the scaling parameters. Adapted from Ref.~\\citenum{cre09}; see the electronic version for colour figures. } \\label{fig:science3} \\end{figure}  Dispensing with tip-tilt completely\\cite{dav08} yields 100\\% sky coverage, and is possible on the VLT because the fast (03\\,Hz) guiding provides a pseudo tip-tilt correction. It opens the door to many important targets that would not otherwise be observable with adaptive optics, with applications ranging from young stellar objects, star clusters, nearby galaxies, and high redshift galaxies. ",
        "conclusions": "We have looked at a number of science cases that have a common goal of understanding the evolution of galaxies and black holes through cosmic time. They employ various techniques of integral field spectroscopy, astrometry, and photometry. Each science case benefits tremendously from -- and is often not possible without -- adaptive optics. Historically, adaptive optics systems have been designed to provide the highest strehl ratios. However, many other issues play a role. The question is then: in order to have the greatest scientific yield, what is the best performance metric?\\vspace{-3mm}  \\begin{description}  \\item{\\em Strehl Ratio}. While a high strehl ratio is always preferable, setting it as a performance metric directly conflicts with extragalactic science, where guide stars or tip-tilt stars are typically faint or sparse.\\vspace{-1mm}  \\item{\\em Spatial Resolution}. Good spatial resolution is mandatory for the science, but reaching the diffraction limit is sometimes not possible due to limitations of the instrumentation (e.g. pixel scale vs field of view) or signal-to-noise. Therefore, although resolution is important in some cases (e.g. photometry of resolved stellar populations), it should not be a primary performance criterion.\\vspace{-1mm}  \\item{\\em Encircled Energy}. This metric embodies the essence of both strehl ratio and resolution, and can be an excellent indicator of performance in terms of the contrast of the high resolution part of the PSF. However, it is often specified within a relatively large diameter. We propose it should be used as a criterion in terms of the fractional energy within the core of the PSF without specifying what the core width should be. This gives some freedom to optimise also other critera, and in some cases even to dispense with a tip-tilt star\\cite{dav08}.\\vspace{-1mm}  \\item{\\em PSF uniformity}. All the science cases require knowledge of the PSF. While the PSF can sometimes be estimated, it remains a major problem when analysing data from current AO systems. An important criterion is therefore that an AO system should enable the PSF to be derived. How this is done is less critical: a uniform PSF across a large field (i.e. MCAO) means it can be measured directly; a parametrization would allow it to be interpolated even if it varies within the field; or ancilliary data might mean there is a way to reconstruct it.\\vspace{-1mm}  \\item{\\em Sky Coverage}. This is perhaps the most important criterion for extragalactic, and in particular for high-z, work. Statistical estimates of the fraction of sources that can be observed at a given galactic latitude do not reflect how widely an AO system can be used. Instead, progress in astrophysics relies heavily on specific targets and specific fields that are intensively studied in many different and complementary ways. An AO system has to be able to provide a good correction on these. The scientific penality paid for insufficient sky coverage is enormous.\\vspace{-1mm}   \\end{description}  \\noindent Our conclusion is that for future AO systems to have a major impact on extragalactic science, they should:\\newline \\indent(i) yield a high fraction of the flux in the PSF core (without specifically optimising strehl or resolution);\\newline \\indent(ii) do so with maximum sky coverage (i.e. be optimised to work at the faint limit); and\\newline \\indent(iii) enable the PSF to be derived (either measured, interpolated, or reconstructed)."
    },
    "1005/1005.2406_arXiv.txt": {
        "abstract": "{  We present measurements of the angular correlation function of galaxies selected from the first field of the H-ATLAS survey. Careful removal of the background from galactic cirrus is essential, and currently dominates the uncertainty in our measurements. For our $250\\ \\mu$m-selected sample we detect no significant clustering, consistent with the expectation that the $250\\ \\mu$m-selected sources are mostly normal galaxies at $z\\lesssim 1$. For our $350\\ \\mu$m and $500\\ \\mu$m-selected samples we detect relatively strong clustering with correlation amplitudes $A$ of 0.2 and 1.2 at $1'$, but with relatively large uncertainties. For samples which preferentially select high redshift galaxies at $z\\sim2-3$ we detect significant strong clustering, leading to an estimate of $r_0 \\sim 7-11 \\,h^{-1}$ Mpc. The slope of our clustering measurements is very steep, $\\delta\\sim2$. The measurements are consistent with the idea that sub-mm sources consist of a low redshift population of normal galaxies and a high redshift population of highly clustered star-bursting galaxies.  }   \\keywords {Submm galaxies --Galaxy clustering } ",
        "introduction": "Submillimetre (sub-mm) selected galaxy samples provide a unique way to trace obscured star formation out to high redshifts (Blain et al. \\cite{blain02}). Models for the evolution of such populations disagree on the nature of the sub-mm sources at high redshifts, with some claiming that they are massive galaxies in the process of forming most of their stellar mass (Granato et al. 2004, Narayanan et al. 2009, Dav\\'e et al. 2010) while others model them as lower mass sources undergoing bursts of star formation with a top heavy IMF (Baugh et al. 2005, Lacey et al. 2009). While evidence on individual sources largely supports a massive halo scenario (Dunne, Eales \\& Edmunds 2003, Swinbank et al. 2008, Michalowski, Hjorth \\& Watson 2010), the best way to measure the statistical halo properties of this population is to measure their clustering. The three-dimensional clustering of sub-mm galaxies provides information about the dark-matter halos that they populate, but direct measurements need distance estimates for each galaxy, which we do not have for our sample. The angular clustering can be measured for flux-limited samples but, in order to compare with models, predictions are required for both the $n(z)$ of flux limited samples and the intrinsic 3-d clustering of the galaxies. The sub-mm colours depend on the source redshift, and so selecting on colour can preferentially select high or low redshift samples (see e.g. Amblard et al. 2010).  Previous attempts to measure the clustering of sub-mm sources have been mostly based on catalogues which cover only very small areas. Despite predictions that sub-mm galaxies should have high spatial clustering, their $n(z)$ is broad and so previous work has had limited success in detecting a significant angular clustering signal (Blain et al. 2004, Scott et al. 2006, Wei\\ss\\ et al. 2010). A more recent approach has been to measure the power spectrum of larger area sub-mm maps from BLAST (Viero et al. \\cite{viero}, Devlin et al. \\cite{blast}). This analysis has found significant evidence for clustering, although at relatively low amplitude.  The {\\it Herschel} ATLAS (H-ATLAS) (Eales et al. \\cite{eales}) will survey over 550 deg$^2$ in 5 wavebands at 100, 160, 250, 350 and $500 \\ \\mu$m. One field covering $\\sim 4^\\circ\\times4^\\circ $ degrees was observed during the science demonstration phase of the mission, and has produced a catalogue of $\\sim 6600$ sources with significance $>5 \\sigma$ above the combined instrumental and confusion noise. This represents roughly 1/30 of the final H-ATLAS data-set.  In this paper we present measurements of the angular correlation function of five flux and colour-selected samples of the H-ATLAS sources. ",
        "conclusions": "\\label{sec:4}  The $S_{250}$ sample has no detectable clustering signal. The simple $S_{350}$ flux limited sample does show fairly significant clustering at scales $\\lesssim 2'$, but the $350+$ sample produces a higher amplitude and more significant detection, as expected given that adding cuts in the other two bands leads to a narrower $n(z)$ by removing low-z galaxies.  The $S_{500}$ sample gives a noisier $w(\\theta)$, but also shows a high amplitude. The colour-selected sample with $S_{500}/S_{250}>0.75$ shows a higher amplitude, and higher noise compared to the $350+$ sample. The power-law fits to all samples give steep slopes $\\delta \\sim 2$.  The increasing amplitude in samples selected at longer wavelengths suggests that clustering is stronger in the higher redshift populations, since selection at longer wavelengths tends to favour higher redshift galaxies. According to some models, these are highly clustered star-burst galaxies which are the ancestors of present day ellipticals (Negrello et al. \\cite{mattia}, Dav\\'e et al. \\cite{dave}). The two samples with additional color cuts have $n(z)$ from photometric redshifts derived by Amblard et al. 2010. which can be used to convert the angular amplitudes to spatial $r_0$.  For both samples we find a range of $r_0 \\sim 7-11h^{-1}$ Mpc, depending on the slope used ($\\delta=2$ or $\\delta=0.8$ respectively). While the slope we measure at scales of a few arcmin is universally steep we cannot be sure with the current data-set what the behaviour will be at larger scales. A larger data set is required to fully address the behaviour of the slope and hence reduce the uncertaintites in $r_0$.  At first sight, the non-detection of clustering in the 250$\\ \\mu$m sample is somewhat surprising. It contains a high fraction of lower redshift galaxies, with $>30$ percent at $z<1$ (Smith et al. in prep). This low-z population is expected to cluster in a similar way to local optical galaxies. Galaxies in the SDSS with magnitudes $21<r^*<22$ have a correlation amplitude of $0.046$ at $1'$ (Connolly et al. 2002). Also most of the 350$\\ \\mu$m sample is a subset of the 250$\\ \\mu$m sample, so their clustering will contribute to the 250$\\ \\mu$m clustering. Assuming they are uncorrelated with the low-z galaxies, they will contribute an amplitude scaled down by the relative density squared, leading to an expected amplitude $\\sim 0.03$. These relatively small amplitudes are consistent with our measurement of $-0.01\\pm 0.07$.  Overall, these results are consistent with the general expectation from models which include highly clustered high redshift and weakly clustered low redshift populations of sub-mm emitting galaxies. Models which postulate that sub-mm sources have a higher mass IMF than normal (Baugh et al. 2005, Lacey et al. 2010) predict a higher clustering strength at lower redshifts, which seems at first glance to be at odds with our result. However, these models do predict far stronger clustering above a threshold in luminosity which is redshift dependent. It is possible that our high-z samples exceed this threshold while the low-z samples do not. It is beyond the scope of the current data set to be able to confirm or rule out these models.  There are few previous observations to compare to directly, and those that are available have rather different selections and redshift distributions.  Magliocchetti et al. (\\cite{manuela}) analysed the distribution of bright 24$\\ \\mu$m sources with faint optical counterparts, split into high ($\\langle z \\rangle \\sim 2$) and low redshift ($\\langle z \\rangle \\sim 0.8$) sub-samples. They found a low amplitude ($A\\sim 0.14 \\pm 0.05 $) at low redshift and a higher amplitude ($A\\sim 0.26\\pm0.1$) at high redshift. This trend of $w(\\theta)$ increasing towards higher z is similar to our observations.  The power spectrum analysis of the BLAST data by Viero et al. (\\cite{viero}) detects clustering on scales $5' < \\theta < 20'$. Once interpreted within the Halo Model formalism, their measurement points to an increase in the spatial clustering of the background sub-mm source population with increasing wavelength and therefore increasing redshift. Again this is consistent with our findings.  LABOCA observations of the extended Chandra Deep Field South have produced a catalogue of 126 sources selected at 850$\\ \\mu$m (Wei\\ss\\ et al \\cite{weiss}). A weak detection of clustering is found on scales $\\theta<2'$. Fixing the slope to be $0.8$, their power-law fit has an amplitude of $0.18\\pm 0.1$ at $1'$. This is similar to the amplitude that we find for the 350$\\ \\mu$m selected sample.  Though all three of these measurements are similar to ours, it is not simple to make a direct comparison because either or both of the flux-limit or pass-bands are different, and so the redshift distributions are not the same.  Given the statistical limitations on the current small field, we leave detailed comparisons to models to a later analysis with more data. However we can say that our measurements appear to be consistent with the prevailing models, where sub-mm sources consist of a low redshift population of normal star-forming disk galaxies that have a spatial correlation length, $r_0 \\sim 4$ Mpc, and a high redshift population of highly clustered star-bursting galaxies $r_0 \\sim 10$Mpc (Negrello et al. \\cite{mattia}, Narayanan et al. \\cite{narayanan}, Dav\\'e et al. \\cite{dave}). The rather steep slope of our measurements for the higher redshift samples is also consistent with the high redshift population forming in compact protoclusters.  The current field is only 1/30 of the area that H-ATLAS will cover, and it has the brightest cirrus background of the planned fields. The final H-ATLAS dataset will have much larger fields with lower cirrus emission, and so will provide a benchmark measurement of source clustering in sub-mm populations."
    },
    "1005/1005.2295_arXiv.txt": {
        "abstract": "{Mixing processes in stars driven by composition gradients as a result of the Rayleigh-Taylor instability are not anticipated. They are supported only by hydrodynamic studies of stellar convection. We find that such mixing occurs below the bottom edge of convection zones in our multidimensional hydrodynamic shell convection models.  It operates at interfaces created by off-center nuclear burning, where less dense gas with higher mean molecular weight is located above denser gas with a lower mean molecular weight.} {We discuss the mixing under various conditions with hydrodynamic convection models based on stellar evolutionary calculations of the core helium flash in a 1.25 \\Msun star, the core carbon flash in a 9.3\\,\\Msun star, and of oxygen burning shell in a star with a mass of 23\\,\\Msun.} {We simulate the hydrodynamic behavior of shell convection during various phases of stellar evolution with the Eulerian hydrodynamics code HERAKLES in two and three spatial dimensions. Initial models for this purpose are obtained by state-of-the-art stellar evolutionary codes, namely GARSTEC, STAREVOL, and TYCHO for the core helium flash, core carbon flash, and oxygen shell burning, respectively. Most of our analysis is performed for two-dimensional hydrodynamic models of shell convection during the core helium flash at its peak covering approximately 250 convective turnover timescales or 1.4 days of stellar evolution.} {The mixing manifests itself in the form of overdense and cold fingers enriched with matter of higher mean molecular weight, originating from density fluctuations at the lower boundary of the convective shell, and ``shooting'' down into the core. The fingers are neither produced by overshooting from the convection zone nor by thermohaline mixing. Instead, they result from the Rayleigh-Taylor instability at the lower convection zone boundary due to a negative mean molecular weight gradient (molecular weight decreasing in the direction of gravity). They do not appear when the mean molecular weight gradient is positive. We call this process Rayleigh-Taylor finger instability mixing or RTFI mixing for short.} {} ",
        "introduction": "\\label{sect:intro}  The hydrodynamic approach to model certain phases of stellar evolution, like \\eg a nuclear core flash, by numerically solving the Euler or Navier-Stokes equations is essentially built upon first principles in physics, and thus (almost) parameter-free. This approach is advantageous compared to (1D) stellar evolutionary calculations when modelling phenomena related to flow instabilities which are inherently multidimensional in nature. One of these instabilities is the Rayleigh-Taylor (RT) instability \\citep{Rayleigh1883, Taylor1949}, caused by the buoyancy force which acts on fluid/gas elements (blobs) of matter that are over-dense or over-light with respect to the matter surrounding them.  One consequence of the RT instability in stellar evolution is the occurrence of convection \\citep{KipWeigert1990, CoxGiuli2008}, whose role for stellar evolution is even nowadays still mainly known from 1D hydrostatic stellar evolutionary calculations. However, convection is a genuine multidimensional phenomenon, \\ie its properties can be reliably inferred only on the basis of multidimensional hydrodynamic simulations. To this end several hydrodynamic studies of shell convection in the cores of stars (hence sometimes called core convection) have been performed in recent years studies, with interesting conclusions \\citep{Deupree1996, Dearborn2006, Herwig2006, MeakinArnett2007, Mocak2009}.  These simulations also considered the effects of composition gradients, but without investigating their influence in much detail, except for the work of \\citet{MeakinArnett2007} highlighting the importance of turbulent entrainment \\citep{Fernando1991} at the boundaries of convection zones.  Composition gradients play an important role during the evolution of stars of all masses. They can instigate double-diffusive instabilities \\citep{Kippenhahn1980, Turner1985} \\ie ``salt-fingering'' or the occurrence of vibrational overstability known as semi-convection \\citep[;and references therein]{Grossman1996} and can trigger gravitational settling \\citep{Stancliffe2008} thereby influence the evolution of a star. Chemical gradients are also likely responsible that a star becomes a red giant \\citep{Stancliffe2009}, and they lead to deep mixing in envelopes of low-mass red giants \\citep{Eggleton2006}.  Stellar layers with a negative mean molecular weight gradient (\\ie the molecular weight $\\mu$ decreases in the direction of gravity) can develop a shell where thermohaline mixing occurs \\citep{Kippenhahn1980}, which is driven by a double-diffusive instability.  Thermohaline mixing can be regarded as a diffusion process where blobs of gas with higher molecular weight sink down, while other material rises and mixes with neighboring gas, eventually providing further rising and sinking elements.  Such a situation is encountered, \\eg during the off-center ignition of helium in the helium core during the core helium flash, where a convective shell forms above the ignition point which is continuously enriched by nuclear ashes with a higher molecular weight (mainly $^{12}$C). This creates a negative mean molecular weight gradient at the base of the convection zone, which should cause thermohaline mixing within a thermal timescale \\citep{Kippenhahn1980}. Thermohaline mixing is supposed to manifest itself in the form of finger-like structures \\citep{Turner1985}, which are indeed observed, \\eg in oceanography \\citet{Williams1974, Gargett1982}. On the other hand, results from experimental chemo-hydrodynamics often reveal instabilities that have a finger-like structure as well, but which are not regarded as thermohaline mixing. Well known in this context \\citep{AnneDeWit2004, AnneDeWit2008} is viscous fingering  where a less viscous fluid is displacing a more viscous one in porous media (Saffamn-Taylor instability), and density fingering driven by a Rayleigh-Taylor instability in systems, where a denser solution resides above a lighter one.  In the following we discuss a mixing process characterized by sinking dense blobs of gas that carry matter of higher mean molecular weight $\\mu$ rapidly and presumably adiabatically downwards (in the direction of gravity) into matter of lower $\\mu$, leaving behind trails which resemble finger-like structures.  This process that has not been anticipated by stellar evolutionary calculations is neither the result of thermohaline mixing, nor viscous fingering, nor density fingering. The mixing process operates below the base of convection zones where the mean molecular weight gradient is negative, \\ie where ``lighter'' matter resides above ``denser'' matter. Hence, it does not occur below convection zones where the mean molecular weight gradient is positive.  As it is caused by a Rayleigh-Taylor instability, we call it Rayleigh-Taylor Finger Instability mixing, or RTFI mixing for short.  Intrinsic properties of the finger-like structures inferred mainly from our hydrodynamic simulations of shell convection during the core helium flash rule out an explanation within the framework of thermohaline mixing as the finger-like structures are not hotter, but colder than the ambient medium. This fact also supports findings of \\citet{Kippenhahn1980}, who estimated based on calculations of \\citet{Thomas1967} for the non-central core helium flash of a 1.3 \\Msun star that the timescale for thermohaline mixing (the time it takes diffusion to double the geometrical thickness of the partially mixed region below the convection zone) is of the order 10$^5\\,$yrs, which is much too slow to be seen in our hydrodynamic simulations. It is neither expected that the timescale will drastically change (decrease) with improved stellar models nor during the flash as the conditions are in general very similar \\citep{SweigertGross1978}.   \\begin{figure} \\includegraphics[width=0.95\\hsize]{formiro.jpg} \\vspace{1.2cm} \\caption{Buoyancy-driven deformation of an interface between a less dense solution of HCl (yellow) residing above a denser solution of NaOH (blue) in an external gravitational field (pointing downwards in the figure). The field of view is 3 by 3\\,cm, and the snapshot is taken few minutes after contact \\citep{Almarcha2010}.} \\label{fig.expfing} \\end{figure}   Contrary to the thermohaline instability, a dynamic mixing process like RTFI mixing is unknown in experimental chemo-hydrodynamics.  It is unexpected as matter is mixed from a less dense layer into a denser one.  However, similarly looking buoyancy-driven instabilities were indeed found in the laboratory at the interface between a solution of lighter hydrochloric acid (HCl) residing above a heavier solution of sodium hydroxide (NaOH) oriented vertically in the gravitional field by \\citet{Zalts2008}. The reaction of these two solutions at their common interface leads to the production of salt (NaCl), and to over-dense (with respect to the horizontal average) fingers of NaCl extending into the NaOH solution as shown in a snapshot from the experiment of \\citet{AnneDeWit2008} or \\citet{Almarcha2010} in Fig.\\,\\ref{fig.expfing}.  Such systems resemble principle similarities with the stellar shell convection models where lighter gas of higher mean molecular weight is residing above denser gas of lower mean molecular weight. Proposed explanations for this kind of laboratory mixing range from amplified local density fluctuations due to salt production, a local density decrease due to a temperature rise up to differences in the mass diffusion coefficients of the reacting species \\citep{Zalts2008}.  Recent work seems to favor the latter one \\citep{Almarcha2010}.  However, mass diffusion in stars is too slow to be responsible for the mixing observed in our simulations (see \\eg \\citet{Michaud2008}).   \\begin{table} \\caption[]{Some characteristic properties of our initial models: total mass $M$, stellar population, metal content $Z$, mass $M_{m}$, outer (inner) radius $R_{m}$, and nuclear energy production rate $L_{m}$, respectively.} \\vspace{0.5cm} \\begin{tabular}{l|lcllll} Model & $M$  & Pop. & $Z$ & $M_{m}$  & $R_{m}$    & $L_{m}$    \\\\ & $[\\Msun]$ &      &     & $[\\Msun]$ & $[10^9\\cm]$ & $[10^9\\Lsun]$ \\\\ \\hline M  & $1.25$ & I & $~~~0.02$ & $0.45$ & $1.2(0.2)$ & $\\sim 1$  \\\\ L  & $9.3 $ & I & $~~~0.02$ & $0.94$ & $1.0(0.3)$  & $\\sim 0.01$ \\\\ O  & $23. $ & I & $\\sim 0.02$ & $2.67$ & $1.0(0.3)$  & $\\sim 1000$  \\\\ \\end{tabular} \\label{imodtab} \\end{table}  \\begin{figure} \\includegraphics[width=0.99\\hsize]{inihefl.pdf} \\vspace{-0.5cm} \\includegraphics[width=0.99\\hsize]{inicfl1.pdf} \\vspace{-0.5cm} \\includegraphics[width=0.99\\hsize]{inioxfl.pdf} \\caption{Temperature $T$ (solid) and mean molecular weight $\\mu$ (dotted) as a function of radius for the initial core helium flash model (a), the core carbon flash model (b), and the oxygen shell burning model (c), respectively . In each of these panels the region of interest below the base of the convection zone is highlighted by the shaded vertical strip. } \\label{fig.inimod} \\end{figure}  The paper is organized as follows. In Sect.\\,\\ref{sect:inidat} we discuss the input models for our simulations. Our hydrodynamic code and the hydrodynamic shell convection models together with the local and temporal properties of the RTFI mixing are described in Sect.\\,\\ref{sect:hydrosim}. A possible theoretical explanation of RTFI mixing is discussed in Sect.\\,\\ref{sect:locstab}, while the breaking of dynamic stability in stars is analyzed in Sect,\\,\\ref{sect:disc}. Finally, a summary of our findings is given in Sect.\\,\\ref{sect:sum}. ",
        "conclusions": "\\label{sect:disc}  Despite the observed existence of dynamic mixing in stars driven by chemical gradients, its inclusion into stellar evolutionary calculations imply non existing observational evidence.  Motivated by mixing driven by the $^3$He inversion found in hydrodynamic simulations of hydrogen layers in low-mass red giants by \\citet{Eggleton2006}, \\citet{Denissenkov2008} discovered that the time which stars spend in the bump while climbing the red-giant branch is significantly prolonged.  These finding suggests that layers with a negative gradient of the mean molecular weight should not become unstable on dynamical timescales. However, we observe just the opposite in our hydrodynamic simulations, where $\\nabla_{sur}-\\nabla_{ad}-\\frac{\\phi}{\\delta} \\nabla_\\mu > 0$ is a very common situation, as demonstrated by Fig.\\,\\ref{fig.nablasplan}. It depicts the stability properties of individual grid zones in the temperature ($\\nabla_{sur} -\\nabla_{ad}$) $vs$ composition gradient $\\frac{\\phi}{\\delta}\\nabla_\\mu$ plane.  Initially, all grid zones above and below the convection zone have a sub-adiabatic temperature gradient ($\\nabla_{sur}-\\nabla_{ad}<0$) and thus are stable, while for grid zones in the convective shell the temperature gradient exceeds the adiabatic one by 10$^{-3}$ rendering them unstable (Fig.\\,\\ref{fig.nablasplan}\\,a). After the onset of convection the situation changes considerably (Fig.\\,\\ref{fig.nablasplan}\\,b). Whereas the temperature gradient remains super-adiabatic in many grid zones inside the convection zone, some of these zones become sub-adiabatic and their composition changes. This can be understood by realizing that convection consists of rising hot plumes (super-adiabatic) of gas enriched by nuclear ash from nuclear burning and sinking cold plumes (sub-adiabatic) mixed with gas from layers further out. Chemical species are mixed quite uniformly throughout the evolved convection zone,\\ie the corresponding grid zones tend to gather at $\\nabla_\\mu \\sim 0$.  On the other hand, some zones from below the base of the convection zone become dynamically unstable as their $\\nabla_{sur}-\\nabla_{ad}-\\frac{\\phi}{\\delta}\\nabla_\\mu$ is positive. These are the zones which participate in RTFI mixing."
    },
    "1005/1005.2540_arXiv.txt": {
        "abstract": "{} {In gravitational lensing the average colors of the images are not identical to the average color of the source. The highly non-linear mapping of gravitational lensing does not preserve the color balance of the source, and this mapping is different for each image.} {The color distortion of the images is illustrated using HST images of the lens SL2SJ02140. It is shown that in this lens the color of the images is variable, reflecting the variable color of the source.} {The average color of the images in SL2SJ02140 are interpreted as a variable amplification of different sources regions with different colors.} {The variation of the average image colors affects the measurements of the photometric redshift of the images. This is especially true for SL2SJ02140 where the color variations due to the non-linear mapping of the lens simulates pseudo redshifts variations.} ",
        "introduction": "Gravitational lensing preserve the surface brightness of the source the color of the local source element. As a consequence it is generally assumed that the average color of the images reflects the color of the source. But this statement is true only if the color of the source is constant. In practice this is rarely the case as illustrated by SL2SJ02140 (Alard 2009) or SL2S02176 (Alard 2010). Since in a gravitational lens the images of the source are highly distorted, the surface occupied by an area of the source of a given color is different in each images. For instance, some source area of a given color are over-amplified and thus over represented. This non-linearity is different in each image, leading to color variations of the images. ",
        "conclusions": "The analysis of SL2SJ02140 demonstrates that the color variations within each image should always be analyzed before making any conclusion about the average image colors. This analysis is absolutely essential before deriving any results about the intrinsic color of the source. One should not forget that the non-linearity of gravitational lensing greatly complicates the interpretation of image colors, and that a straightforward interpretation of the average image colors may be wrong. For instance Verdugo {\\it etal.} 2010 found that using the average color of the images the pseudo redshift of the source was between 1.0 and 1.8, which illustrates the amplitude of the errors on the redshift that can be made when the non-linearity of the lens mapping is ignored."
    },
    "1005/1005.0883_arXiv.txt": {
        "abstract": "A detection of the level of non-Gaussianity in the CMB data is essential to discriminate among inflationary models and also to test alternative primordial scenarios. However, the extraction of primordial non-Gaussianity is a difficult endeavor since several effects of non-primordial nature can produce non-Gaussianity. On the other hand, different statistical tools can in principle provide information about distinct forms of non-Gaussianity. Thus, any single statistical estimator cannot be sensitive to all possible forms of non-Gaussianity. In this context, to shed some light in the potential sources of deviation from Gaussianity in CMB data it is important to use different statistical indicators. In a recent paper we proposed two new large-angle non-Gaussianity indicators which provide measures of the departure from Gaussianity on large angular scales. We used these indicators to carry out analyses of non-Gaussianity of the bands and of the foreground-reduced WMAP maps with and without the \\emph{KQ75} mask. Here we briefly review the formulation of the non-Gaussianity indicators, and discuss the analyses made by using our indicators. ",
        "introduction": "A detection of the level of primordial non-Gaussianity in the CMB data is crucial to discriminate inflationary models and also to test alternative scenarios for the physics of the early universe. Clearly the study of detectable non-Gaussianities in the WMAP data ought to take into account that they may have non-cosmological origins as, for example, unsubtracted foreground contamination, unconsidered point sources emission and systematic errors. % Deviation from Gaussianity may also have a cosmic topology origin (see, e.g., the review Refs.~\\refcite{CosmTopReviews} and  related Refs.~\\refcite{TopDetec}). If, on the one hand, different statistical tools can in principle provide information about distinct forms of non-Gaussianity, on the other hand one does not expect that a single statistical estimator can be sensitive to all possible forms of non-Gaussianity in CMB data. In view of this, a great deal of effort has recently gone into verifying the existence of non-Gaussianity by employing several statistical estimators (an incomplete list of references is given, e.g., in Refs.~\\refcite{Some_non-Gauss-refs,Bernui-Reboucas2009a,Bernui-Reboucas2009b} and references therein).  Recently we have proposed\\cite{Bernui-Reboucas2009a} two new large-angle non-Gaussianity indicators, based on skewness and kurtosis of large-angle patches of CMB maps, which provide measures of the departure from Gaussianity on large angular scales. We used these indicators to search for the large-angle deviation from Gaussianity in both band and foreground-reduced maps\\cite{Bernui-Reboucas2009b} with and without a \\emph{KQ75} mask (see also the related Refs.~\\refcite{Related}). Here we briefly summarize the main results of Refs.~\\refcite{Bernui-Reboucas2009a} and~\\refcite{Bernui-Reboucas2009b}. %   \\vspace{-2mm} ",
        "conclusions": "A constructive way of formulating our non-Gaussianity indicators $S$ and $K$  from CMB data is through the following steps: \\begin{alphlist} \\item Take a finite set of points $\\{j=1, \\ldots ,N_{\\rm c}\\}$ homogeneously distributed on the CMB  celestial sphere $S^2$ as the centers of spherical caps of a given aperture $\\gamma$; and calculate for each cap $j$ the skewness ($S_j$) and kurtosis ($K_j$) by using that \\begin{equation} S_j   \\equiv  \\frac{1}{N_{\\rm p} \\,\\sigma^3_{\\!j} } \\sum_{i=1}^{N_{\\rm p}} \\left(\\, T_i\\, - \\overline{T_j} \\,\\right)^3 \\quad \\mbox{and} \\quad K_j   \\equiv  \\frac{1}{N_{\\rm p} \\,\\sigma^4_{\\!j} } \\sum_{i=1}^{N_{\\rm p}} \\left(\\,  T_i\\, - \\overline{T_j} \\,\\right)^4 - 3 \\,, \\end{equation} where $N_{\\rm p}$ is the number of pixels in the $j^{\\,\\rm{th}}$ cap, $T_i$ is the temperature at the $i^{\\,\\rm{th}}$ pixel, $\\overline{T_j}$ is the CMB mean temperature of the $j^{\\,\\rm{th}}$ cap, and $\\sigma$ is the standard deviation. The numbers $S_j$ and $K_j$ obtained in this way for each cap can be seen as a measure of non-Gaussianity in the direction of the center $(\\theta_j, \\phi_j)$ of the $j^{\\,\\rm{th}}$ cap. \\item Patching together the $S_j$ and $K_j$ values for each spherical cap, one obtains two discrete functions $S = S(\\theta,\\phi)$ and $K = K(\\theta,\\phi)$ defined on the sphere $S^2$, which can be used as statistical indicators to measure the deviation from Gaussianity as a function of the angular coordinates $(\\theta,\\phi)$. The Mollweide projection of skewness and kurtosis functions $S = S(\\theta,\\phi)$ and $K = K(\\theta,\\phi)$  are nothing but skewness and kurtosis maps (hereafter $S-$map and $K-$map). \\end{alphlist}  Clearly, the discrete functions $S = S(\\theta,\\phi)$ and $K = K(\\theta,\\phi)$ can be expanded into their spherical harmonics in order to determine their power spectra $S_{\\ell}$ and $K_{\\ell}$. Thus, for example, for the skewness one has $S (\\theta,\\phi) = \\sum_{\\ell=0}^\\infty \\sum_{m=-\\ell}^{\\ell} b_{\\ell m} \\,Y_{\\ell m} (\\theta,\\phi)$ and $S_{\\ell} = (2\\ell+1)^{-1}\\sum_m |b_{\\ell m}|^2$. % Similar expressions obviously  hold for the kurtosis $K = K(\\theta,\\phi)$.  In the remainder of this work we shall report the results of our Gaussianity analysis performed with  $S = S(\\theta,\\phi)$ and $K = K(\\theta,\\phi)$ indicators calculated from single frequency and foreground reduced maps with and without  a \\emph{KQ75} mask. To minimize the statistical noise, in the calculations of $S-$map and $K-$map  from the input maps, we have scanned the celestial sphere with spherical caps of aperture  $\\gamma = 90^{\\circ}$, centered at $12\\,288$ points homogeneously distributed on the two-sphere.  For the sake of brevity we do note show examples of $S$ and $K$ maps, which provide only \\emph{qualitative} information on large-angle deviation from Gaussianity (see figures in Refs.~\\refcite{Some_non-Gauss-refs,% Bernui-Reboucas2009a,Bernui-Reboucas2009b} for examples). To obtain \\emph{quantitative} information about the large angular scale (low $\\ell$) distributions for the non-Gaussianity $S$ and $K$ maps obtained from the  CMB input maps used, we have calculated the (low $\\ell$) power spectra $S_{\\ell}$ and $K_{\\ell}$ for $S$ and $K$ maps. The deviation from Gaussianity and the statistical significance were estimated by comparing these power spectra with the corresponding averaged power spectra $\\overline{S}_{\\ell}$ and $\\overline{K}_{\\ell}$ calculated from $S$ and $K$ maps obtained by averaging over $1\\,000$ Monte-Carlo-generated statistically Gaussian CMB maps. To have an overall assessment of low $\\ell$ power spectra $S_\\ell$ and $K_\\ell$ calculated from each CMB input, we have performed a $\\chi^2$ test to find out the goodness of fit for $S_{\\ell}$ and $K_{\\ell}$ multipole values as compared to the expected multipole values from the MC Gaussian maps. In this way, we obtained one number for each map that collectively ('globally') quantifies the deviation from Gaussianity. For the power spectra $S_\\ell$ calculated from $S-$maps obtained from the five-years maps with a \\emph{KQ75} mask we found that the ratio $\\chi^2/\\,\\text{dof}\\,$ (dof stands for degree of freedom) from the K, Ka, Q, V, and W  maps  are given, respectively, by $21.5$, $4.9$, $6.0$, $5.2$, and $3.9$,  while for the kurtosis power spectra $K_{\\ell}$ of these maps the values of $\\chi^2/\\text{dof}$ are, respectively, $35\\,652$, $135$, $0.5$, $6.4$, and  $5.6$. Clearly a good fit occurs when $\\chi^2/\\text{dof}\\,\\sim 1 $. Moreover, the greater are the $\\chi^2/\\text{dof}\\,$  values the smaller are the $\\chi^2$ probabilities, that is the probability that the multipole values $S_{\\ell}$ and $K_{\\ell}$ and the expected MC multipole values agree. For $S_{\\ell}$ and $K_{\\ell}$ obtained from the full-sky foreground-reduced five years ILC input maps we found that $\\chi^2/\\text{dof}\\,$ are, respectively,  $35.7$  and $2\\,368$. These results reduce to  $1.2$  and $0.4$ when the \\emph{KQ75} mask is employed. In brief, our analyses show that the unmasked band maps are significantly non-Gaussian but the deviation from Gaussianity is substantially reduced to a level compatible with Gaussianity for Q and V maps, whereas the full-sky foreground reduced five years ILC mask is  again significantly non-Gaussian but the level of non-Gaussianity drops to a level that is consistent with Gaussianity when the \\emph{KQ75} mask is used.  \\vspace{-3mm}"
    },
    "1005/1005.5004_arXiv.txt": {
        "abstract": "{In the recent years, Stellar Mass Functions of both star forming and quiescent galaxies have been observed at different redshifts in various fields. In addition, Star Formation Rate distributions (e.g. in the form of far infrared luminosity functions) were also obtained. Taken together, they offer complementary pieces of information concerning the evolution of galaxies.} {We attempt in this paper to check the consistency of the observed stellar mass functions, SFR functions and the cosmic star formation rate density with simple backward evolutionary models.} {Starting from observed Stellar Mass Functions for star-forming galaxies, we use backwards models to predict the evolution of a number of quantities, such as the SFR function, the cosmic SFR  density and the Velocity Function. The velocity being a parameter attached to a galaxy during its history (contrary to the stellar mass), this approach allows us to quantify the number density evolution of galaxies of a given velocity, e.g. of the Milky Way siblings. } {Observations suggest that the Stellar Mass Function of star forming galaxies is constant between redshift 0 and 1. In order to reproduce this result, we must quench star formation in a number of star forming galaxies. The Stellar Mass Function of these ``quenched'' galaxies is consistent with available data concerning the increase in the population of quiescent galaxies in the same redshift interval. The Stellar Mass Function of quiescent galaxies is then mainly determined by the distribution of active galaxies that must stop star formation, with a modest mass redistribution during mergers. The cosmic SFR density, and the evolution of the SFR functions are relatively well recovered, although they provide some clue for a small evolution of the Stellar Mass Function of star forming galaxies at the lowest redshifts. We thus consider that we have obtained in a simple way a relatively consistent picture of the evolution of galaxies at intermediate redshifts.  We note that if this picture is correct, 50\\% of the Milky-Way sisters (galaxies with the same velocity as our Galaxy, i.e. 220 km/s) have quenched their star formation since redshift 1 (and an even larger fraction for larger velocities). We discuss the processes that might be responsible for this transformation.} {} ",
        "introduction": "In the last few years, the analysis of many deep fields together with improvements in techniques such as  SED-fitting have enabled many studies of the Stellar Mass Function (SMF) of galaxies at both low and high redshifts (e.g. Bell et al. 2003, 2007, Borch et al. 2006, Bundy et al. 2005, 2006, Arnouts et al. 2007, Vergani et al. 2008, Perez-Gonzalez et al. 2008, Pozzetti et al. 2009, Ilbert et al. 2009). One of the somewhat surprising result is that it is generally found that the SMF of star forming galaxies (selected on the basis of color, morphology or spectroscopy) has not evolved much between redshift about unity and the present time (e.g. Bell et al. 2007, Cowie et al. 2008, Vergani et al. 2008, Pozzetti et al. 2009). The analysis of the evolution of the Luminosity Function had already suggested that the stellar mass of blue galaxies is constant, while the one of red galaxies increases with time (Faber et al. 2007).   Since star formation is undergoing in such galaxies, the stellar mass in each of them is bound to increase.  This maintained status quo of the SMF thus indicates that a number of massive star forming galaxies must evolve towards early type, or at least quiescent galaxies (through quenching of their star formation), as noted in \\cite{bell07} and \\cite{arnouts07}, while they are replaced by lower mass galaxies gaining mass with time (see also Pozzetti et al. 2009). This type of scenario was also advanced in Brown et al. (2007) to explain the doubling of the stellar mass in red galaxies over the past 8 Gyrs. Such a quenching is also constrained by the evolution of the co-moving Star Formation Rate (SFR) density, i.e. the SFR averaged over representative volumes of the universe, dropping by a factor 6 to 10 between redshift 1 and 0 (see the compilations of Wilkins et al.  2008 and Hopkins and Beacom 2006).    It is reasonable to think that even at ``intermediate'' redshifts (defined broadly as $0<z<1$), star formation takes place in ``normal'' star forming galaxies, such as the precursors to present day disk-galaxies and the Milky Way. Smooth SFR histories are known to be realistic and capture enough of the physics of such galaxies (see e.g. models in Boissier \\& Prantzos 1999, 2000). This is especially suggested by the facts that we see star formation at intermediate redshift in rotating disks that should have formed through smooth accretion (e.g. Shapiro et al. 2008) and that most of the galaxies forming the bulk of the infra-red emission at redshift unity look like undisturbed disk galaxies (e.g. Bell 2005, Melbourne et al. 2005, Lotz et al. 2008). The tightness of the SFR-Stellar Mass relationship also suggests that the evolution is mostly driven by smooth evolution and not starbursts (Noeske 2009, Noeske et al.2007). The kinematics and morphology are more disturbed at high redshift than in the local universe but this fact is still consistent with minor merger and gas accretion (e.g. Ziegler et al. 2009) although this is still debated (e.g. Neichel et al. 2008). Finally, it has also been shown that models of the chemical and spectro-photometric evolution of the Milky Way and nearby disk galaxies, when rolled backwards, fit in a very reasonable way the properties of redshift $\\sim$ 0.7 star forming galaxies (Buat et al. 2008).  In this paper, we propose to combine such models with the information from the SMF in order to compute the evolution of the ``Velocity Function'', with the advantage that the velocity tags a galaxy, while the stellar mass evolves with redshift when star formation takes place (e.g. Noeske 2009  mentions the intrinsic difficulty in comparing galaxies of a given  stellar mass at different redshifts). This allows us to check the global consistency of the SMF, the SFR of individual galaxies, the cosmic SFR density and simple models of galactic evolution. It is then possible to derive the fraction of galaxies that were similar to the Milky Way (the ``Milky Way siblings'') and that must have quenched their star formation during the second half of the history of the universe. We acknowledge that this approach is simplistic but we think it is worth performing it since it allows us to step back and have a global look at all the important data of the problem (Stellar Mass Functions of quiescent and star forming galaxies, SFR functions, cosmic SFR  density) in a simple, educational framework, while many studies focus only on some of these aspects (with however interesting exceptions as e.g. Hopkins et al. 2009 or Bell et al. 2007).   In section \\ref{secModels}, we present the ``backward'' models of the evolution of star forming galaxies that we use. In section \\ref{secVel}, we compute the velocity function and its evolution assuming a constant SMF, as suggested by the observations. In section \\ref{secEvo}, the consequences of our simple assumptions are discussed. In section \\ref{SecConclu}, a summary of our work is proposed. ",
        "conclusions": "\\label{SecConclu} This paper presents an original and educational approach to work with Stellar Mass Function of star forming galaxies. The main point consists in combining observed Stellar Mass Functions with simple backward evolutionary models (assuming a smooth evolution driven by star formation in a galactic disk fuelled by cold accretion) to obtain modeled velocity functions at various redshifts. The big advantage of the velocity is that it is a model parameter attached to a galaxy, contrary to stellar mass, which evolves with star formation. In the framework of such models, for a given velocity, the Star Formation History and the Stellar Mass history are known. Galaxies with large velocities are the precursor of today's massive star forming galaxies (in the absence of quenching), with stars having formed at an early epoch. Galaxies with low velocities have been formed on average at later times. These histories are not adhoc models for this paper, but have been shown to be consistent with a large number of properties of nearby star-forming galaxies.  Combining the Velocity Functions and the models, it is possible to compute the model predictions for a large number of quantities and distributions at any redshift (e.g. SFR functions and cosmic SFR density).  At the current time, the empirical  stellar mass functions are still rather uncertain (see the scatter between various studies). Assuming that, to the first order, the stellar mass function is constant for star forming galaxies at redshift lower than 1 (in accordance with the data), we obtained several interesting results:   \\begin{enumerate}  \\item A globally consistent picture is obtained in which a large fraction of star forming galaxies quench their activity of star formation between redshift 0 and 1, building up the red sequence. The shape of the quiescent galaxies stellar mass function is recovered above $\\sim$ 3 10$^{10}$ \\msun.  Mergers can only alter it slightly by redistributing the mass among galaxies. At most, their main effect could be seen only on the massive end of the stellar Mass Function.  \\item The evolution of the cosmic SFR density falls within the compilation of \\cite{wilkins08}. However, our decrease towards low redshifts is smaller than in the observations. The discrepency is especially due to our redshift 0 point. This may be related to a small inconsistency between the stellar mass function, SFR function, and stellar mass-SFR relationship of the local universe star forming galaxies.  \\item The fraction of galaxies having quenched their star formation is high, in rough agreement with measurements.  For example, we conclude that half of the Milky Way siblings (galaxies with circular velocities of 220 km/s) have quenched their star formation since redshift 1, and are then no more star-forming disk such as the Galaxy.  This fraction is probably a lower limit since the model cosmic SFR density presents itself a relatively weak evolution with respect to the observed one.  \\item The process responsible for this change is not clearly identified today. What is clear is that it is probably not related to cluster physics or to major mergers (which affect galaxies in too small numbers). Groups are  our prime suspects because they are numerous enough to have a significant effect on the cosmic evolutions and several studies indicate that transformations are indeed taking place in them.  \\end{enumerate}  In the future, we hope that wider, deeper fields and improvement in the analysis of stellar mass functions will allow us to use the method exposed in this paper but on safer grounds. This will allow us to improve our models of galaxy evolution and to get a more complete picture of the  cosmic evolution of star-forming  galaxies, and their transformation into quiescent ones."
    },
    "1005/1005.2412_arXiv.txt": {
        "abstract": "{ We present colour-colour diagrams of detected sources in the {\\it Herschel}-ATLAS Science Demonstration Field from $100$ to $500\\,\\mu$m using both PACS and SPIRE. We fit isothermal modified black bodies to the spectral energy distribution (SED) to extract the dust temperature of sources with counterparts in Galaxy And Mass Assembly (GAMA) or SDSS surveys with either a spectroscopic or a photometric redshift. For a subsample of 330 sources detected in at least three FIR bands with a significance greater than 3$\\sigma$, we find an average dust temperature of $(28 \\pm 8)$K. For sources with no known redshift, we populate the colour-colour diagram with a large number of SEDs generated with a broad range of dust temperatures and emissivity parameters, and compare to colours of observed sources to establish the redshift distribution of this sample. For another subsample of 1686 sources with fluxes above 35 mJy at $350\\,\\mu$m and detected at 250 and $500\\,\\mu$m with a significance greater than 3$\\sigma$, we find an average redshift of $2.2 \\pm 0.6$. } ",
        "introduction": "The total intensity of the extragalactic background light at far-infrared (FIR) wavelengths has been measured with absolute photometry (Puget et al. 1996; Fixsen et al.  1998; Dwek et al. 1998). Deep surveys are now starting to resolve the cosmic FIR background into discrete sources with the fraction resolved varying with wavelength (see reviews in Blain et al. 2002; Hauser $\\&$ Dwek 2001;  Lagache et al. 2005, Dye et al. 2009). Despite some successes, there are still large uncertainties about the nature and evolution of the submillimetre (submm) galaxy population, including their redshift distribution.  The {\\it Herschel}-Astrophysical Terahertz Large Area Survey (H-ATLAS)  is an open-time key program of the {\\it Herschel} Observatory (Pilbratt et al. 2010) that will survey roughly 550 deg$^2$ of sky over 600 hours of observations (Eales et al. 2010). To cover the widest possible area, H-ATLAS  uses the maximum possible scan rate for the telescope at 60 arcsec\\,sec$^{-1}$ in the parallel mode of PACS (Poglitsch et al. 2010) and SPIRE  (Griffin et al. 2010). The observations cover five photometric bands with $100$ and $160\\,\\mu$m data from  PACS, and $250$, $350$, and $500\\,\\mu$m data from SPIRE. The H-ATLAS fields include 1 field close to the northern Galactic pole (150 deg$^2$), 3 fields (each 36 deg$^2$) coinciding with the GAMA redshift survey fields (Driver et al. 2009), and 2 fields (total area of 250 deg$^2$) near the southern Galactic pole.  During the Science Demonstration Phase (SDP), H-ATLAS observed 14.4 deg$^2$ in the GAMA-9 hour field near the ecliptic plane to a 5$\\sigma$ depths of 35-90 mJy.  In this {\\it Letter} we present the colour-colour diagrams derived from the five bands from $100$ to $500\\,\\mu$m for the H-ATLAS sources that are cross-identified in at least 3 of the SPIRE and PACS bands. We then discuss the dust temperature and redshift distributions based on certain assumptions about the spectral energy distribution (SED) and subject to selection effects associated with the sample selection.   \\begin{figure}[h!] \\centering \\includegraphics[width=8cm]{14586fg1a.ps} \\includegraphics[width=8cm]{14586fg1b.ps} \\includegraphics[width=8cm]{14586fg1c.ps} \\caption{Colour-colour diagram of SPIRE and PACS sources. The coloured background indicates the average redshift in these colour-colour spaces of our 10$^6$ randomly generated model SEDs. In (a) points plotted represent a sample of 1686 galaxies detected in all three SPIRE bands with a significance greater than 5 $\\sigma$ at $350\\,\\mu$m ($S>35$mJy) and greater than 3$\\sigma$ at $250$ and $500\\,\\mu$m (corresponding to flux limits of roughly 21 and 27 mJy, respectively). In (b) we show a sample of 402 galaxies detected at $160$, $250$, and $350\\,\\mu$m and in (c) we show a sample of 158 galaxies detected at $100$, $160$, and $250\\,\\mu$m. As in (a), in both (b) and (c) we also select sources by imposing a 5$\\sigma$ cut in the central wavelength of the colour plot and greater than 3$\\sigma$ at other wavelengths. } \\label{colcol} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=9cm]{14586fg2.ps} \\caption{Dust temperature as a function of redshift for 330 H-ATLAS sources reliably identified with GAMA and SDSS DR-7 galaxies having a known spectroscopic (black filled triangle) or a photometric (black empty triangle) redshift. The selected sources are $>5\\sigma$ detections in at least one band of  either PACS or SPIRE with $3 \\sigma$ flux measurements in at least 2 more bands. The isothermal spectral fits assume $\\beta=1.5$. For other data plotted, see Section~3.} \\label{Tvsz} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=9cm]{14586fg3.ps} \\caption{Dust temperature as a function of the FIR luminosity for 330 H-ATLAS sources reliably identified with GAMA and SDSS DR-7 galaxies with a known spectroscopic (black filled triangle) or a photometric (black empty triangle) redshift. The FIR luminosity is obtained by fitting the PACS and SPIRE measurements with a modified black body (parameters are the dust temperature and the luminosity density, $\\beta$ is fixed to 1.5) and this model SED is integrated between $8$ and $1100\\,\\mu$m. We also indicate BLAST and H-ATLAS best-fit $T_d-\\log (L_{FIR})$ relation with  a dashed blue line (from Dye et al. 2009) and a dashed red line, respectively. } \\label{TvsL} \\end{figure} ",
        "conclusions": "We have discussed the spectral energy distribution of sub-mm galaxies in a 14 deg.$^2$, subset of the full H-ATLAS survey. We have used the colours in 5 bands to measure the dust temperature for the sub-sample of sources with known redshifts. We have also derived a qualitative estimate of the statistical redshift distribution of bright $350\\,\\mu$m selected galaxies assuming broad priors in temperature and the dust emissivity. We found that the average dust temperature of H-ATLAS sources is $28 \\pm 8$K, while $350\\,\\mu$m sources with fluxes above 35 mJy   have an average redshift of $2.2 \\pm 0.6$."
    },
    "1005/1005.4234_arXiv.txt": {
        "abstract": "We construct a simple, spherical blastwave model to estimate the pressure   structure of the intergalactic medium surrounding hyper-starburst galaxies, and argue that the effects of interaction with star-forming galaxy winds may   be approximated at early times by an adiabatically expanding, self-similar   `bubble' as described by \\citet{weaveretal77} and \\citet{ostrikermckee88}. This model is used to make observational predictions for the thermal Sunyaev-Zel'dovich effect in the shocked bubble plasma.  Radiative cooling losses are explored, and it is found that bremsstrahlung will limit the epoch of adiabatic expansion to $10^7$--$10^8$ years: comparable to total hyper-starburst lifetimes.  Prospects for making a first Sunyaev-Zel'dovich detection of galaxy wind bubbles using the Atacama Large Millimeter Array are examined for a number of active hyper-starburst sources in the literature. ",
        "introduction": "After the epoch of reionization, free electrons in the intergalactic medium (IGM) interact with photons in the cosmic microwave background radiation (CMB) via Thompson scattering.  Known as the Sunyaev-Zel'dovich (SZ) effect (e.g.\\ \\citealp{sunyaevzeldovich70,sunyaevzeldovich72}; see also \\citealp{carlstrometal02} for a review), the result is a characteristic spectral distortion of the CMB.  The effect can be split into two contributions: that due to the thermal motion of free electrons which causes an overall increase in the apparent temperature of CMB photons (called the thermal SZ or tSZ effect), and that due bulk kinetic motion (the kSZ effect; see \\citealp{carlstrometal02}).  The tSZ effect, in particular, shows promise as a means of observing the hot plasma surrounding galaxy clusters and other peaks in the matter power spectrum, and recent years have seen a steady growth in both theoretical predictions (e.g.\\ \\citealp{moodleyetal09}; \\citealp*{scannapiecoetal08}; \\citealp{roncarellietal07}) and actual measurements for massive galaxy clusters (e.g.\\ \\citealp{halversonetal09,basuetal10,nordetal09}). It is interesting to consider what other astrophysical objects may be observable in the near future with the tSZ effect (e.g.\\ \\citealp{yamadaetal10}).  In recent years an increasing number of high redshift `hyper-starbursts' (which we define as objects with an estimated star formation rate of $\\dot{M}_{*} \\gtrsim 10^3$ M$_{\\odot}$yr$^{-1}$) have been reported in the literature, often but not always identified with submillimetre galaxies (SMGs), and are believed to represent a significant fraction of the total star formation at these epochs (e.g.\\ \\citealp{blainetal02,solomonvandenbout05,wangetal08,caseyetal09,martinezetal09,riechersetal09,waggetal09}). Other recent results suggest that the galaxy wind outflows frequently observed in star-forming galaxy spectra (see, e.g., \\citealp*{veilleuxetal05} for a review) are ubiquitous wherever there is star formation activity to drive them \\citep{grimesetal09,rubinetal10,weineretal09,krugetal10}, with wind masses and velocities increasing with $\\dot{M}_{*}$. The winds associated with hyper-starbursts at high redshift will be extremely energetic, leading to violent interaction with the plasma of the surrounding intergalactic medium (IGM).  The following question naturally arises: can we detect the tSZ due to the wind-IGM interaction for these events?  A similar question has been posed for active galaxies (e.g.\\ \\citealp*{yamadaetal99}; \\citealp{plataniaetal02,chatterjeekosowsky07,scannapiecoetal08,chatterjeeetal08,yamadaetal10}) and for the combined signal due to star-forming winds (e.g.\\ \\citealp*{majumdaretal01,babichloeb07}), but individual hyper-starburst objects now also raise interesting possibilities. The temperature increase of CMB photons due to the tSZ effect is redshift independent (e.g.\\ \\citealp{carlstrometal02}), which potentially opens a new observational window on extreme star formation at high redshift.  We examine this question by constructing a simplified adiabatic model of the wind-IGM interaction (Section \\ref{sect:bubble}), and then calculate the tSZ effect according to such a model for a variety of possible hyper-starburst targets in the literature (Sections \\ref{sect:sz} \\& \\ref{sect:targets}) along with observational prospects for the upcoming Atacama Large Millimeter Array (ALMA: see \\citealp*{brownetal04}) facility.  In Sections \\ref{sect:cooling} we test the adiabatic expansion assumption by calculating radiative cooling timescales for the hot plasma, and end in Section \\ref{sect:conc} with a discussion and conclusions. ",
        "conclusions": "\\label{sect:conc} We have presented an adiabatic model of wind-driven bubble shockwaves around hyper-starburst galaxies and made predictions for the strength of the tSZ effect due to such objects.  Assuming this model to be a reasonable approximation we have shown that such bubbles are detectable using ALMA, albeit requiring integration times of multiple days (in two or, preferably, three bands). We have shown that bremsstrahlung will dominate the radiative cooling of the shocked gas, and found that this will cause the evolution of the shock to become significantly non-adiabatic after a time comparable to the typical total duration of hyper-starbursts themselves.  In order to construct the model a number of simplifications have been made. Among these, the assumption of spherical symmetry and the homogeneous density $\\rho_d$ are least likely to be valid.  The former assumption may be thrown into question for systems with active galactic nuclei, as are often seen for hyper-starbursts, or with star formation occurring in an extended disc rather than a central core: \\citet{dallavecchiaschaye08} found a bi-polar pressure distribution when modelling outflows in massive disc-like formation, once winds were fully coupled to the gas dynamics of the system. A bi-polar outflow would alter the pressure distribution of the shocked gas and might reduce the strength of the tSZ effect along certain lines of sight.  However, the star typical star formation rates in the massive systems simulated by these authors were small when compared to the figures in Table \\ref{tab:targets}, and it might be expected that the greater mechanical energy input from hyper-starbursts will alter the wind-IGM interaction significantly. \\citet{dallavecchiaschaye08} found interesting qualitative differences in the behaviour of outflows for their simulated disc systems of total halo mass 10$^{10}h^{-1}$M$_{\\odot}$ and 10$^{12}h^{-1}$M$_{\\odot}$ that produced early star formation rates of $\\simeq 0.05$~M$_{\\odot} \\textrm{yr}^{-1}$ and $\\simeq 10 $~M$_{\\odot} \\textrm{yr}^{-1}$ respectively.  An increase of a further two orders of magnitude in star formation rate may cause sufficient increases in wind ram pressure as to qualitatively change the outflow behaviour once more. In a related study \\citet{chatterjeekosowsky07} and \\citet{chatterjeeetal08} found that quasar-driven bubbles in simulated halos provided a reasonable match to a simple spherical model, although again at lower energies and typically after longer timescales than we consider here.  Unfortunately there is insufficient knowledge and modelling of the wind-interstellar medium interaction in these powerful systems to do more than speculate.  The typical morphologies and environments of the star-forming gas regions prior to the onset of strong winds are also poorly known for these hyper-starbursts, which further complicates accurate simulation.  More detailed theoretical modelling combined with future precision observations of the tSZ might help shed light on exactly these questions.  Another simplification, the assumption of a homogeneous IGM, may misrepresent the evolution of the bubble with time: for a more realistic density profile the bubble will initially expand more slowly through the dense gas at the halo centre, and then more rapidly at later times as $\\rho_d(R_2)$ decreases. This is an area in which the modelling could be improved, particularly if aided by more detailed knowledge of the gas density around hyper-starbursts.  A related simplification is the model's neglect of the gravitational deceleration of the shockwave due to the host halo potential.  This effect will generally represent a small correction to the dynamics of adiabatic expansion and radiative losses, which can be seen by comparing the galaxy wind velocities to the escape velocities of most galactic haloes (entrained clouds, however, might impart a velocity kick that prevents their escape). These effects should be considered in more accurate modelling, perhaps using a simplified model of a galaxy halo potential and mass entrainment (e.g.\\ \\citealp*{murrayetal05}).  Important uncertainties also arise due to the assumed values of the starburst age $t$, and the fiducial parameters $v_{\\mw}$ and $\\dot{M}_*$ (or equivalently $f_{\\textrm{w}}$).  It should be emphasized that the fiducial predictions may be adapted as appropriate using the simple scaling relations presented, which may be necessary as we gain a clearer picture of hyper-starburst physics. As an example, using $y \\propto f_{\\mw}^{3/5}$ from equation \\eqref{eq:ypeak} and adopting the value $f_{\\mw} \\simeq 0.4$ of \\citet{dallavecchiaschaye08} reduces the tSZ signals listed in Table \\ref{tab:targets} by $\\simeq 20\\%$. However, a drastic reduction in $f_{\\mw}$ (appropriate if the kinetic energy from hyper-starburst supernovae is mostly lost radiatively) might yet reduce $y$ to the cosmological confusion limit calculated by \\citet{roncarellietal07}. Conversely, observations of the tSZ around hyper-starbursts could in this way, even if not detected, provide interesting constraints on $f_{\\textrm{w}}$ and thus supernova wind feedback in these objects, which may contribute much of the total star-formation at high redshift \\citep{blainetal02}.  They might also constrain the feedback energy injected by the active galactic nuclei commonly hosted in hyper-starburst systems (e.g.\\ \\citealp{chatterjeeetal08}), although more careful modelling will be required to distinguish the two signals.   To conclude, wind-driven bubbles around hyper-starburst objects may be detectable via the tSZ in the near future using submillimetre instruments such as ALMA or its successors. Despite many simplifying assumptions in the model, the enormous mechanical energy of the wind-IGM interaction will ensure that a system approximating the simple, self-similar blastwave must persist while radiative losses are small.  According to such a model, such objects as those listed in Table \\ref{tab:targets} (with estimated $y_{\\textrm{peak}} > 2 \\times 10^{-5}$) might be securely detected with ALMA within days. The number of known hyper-starbursts in the Southern Hemisphere must be expected to increase with time, and one of these may present an excellent opportunity for a first bubble detection with ALMA."
    },
    "1005/1005.3414_arXiv.txt": {
        "abstract": "{ The paper deals with several properties of the Oort cloud of comets. Sun, Galaxy (and Jupiter) gravitationally act on the comets. New physical model of galactic tide is considered. The main results can be summarized as follows: \\\\ 1. Mass of the Oort cloud of comets is less than 1 mass of the Earth ($M_{E}$), probably not greater than 1/2 $M_{E}$.  \\\\ 2. Theoretical number of long-period comets with perihelion distance $q$ $<$ 5 AU is about 50-times greater than the conventional approach yields. Gravity of Jupiter was taken into account in finding this result. \\\\ 3. Semi-major axis $a$ and period of oscillations $P$ of eccentricity (and other orbital elements) are related as $a^{3}$ $P$ $=$ 1 in natural units for a moving Solar System in the Galaxy. The natural unit for time is the orbital period of the Solar System revolution around the galactic center and the natural unit for measuring the semi-major axis is its maximum value for the half-radius of the Solar System corresponding to the half-radius of the Oort cloud. The relation holds for the cases when comets approach the inner part of the Solar System, e.g., perihelion distances are less than $\\approx$ 100 AU. \\\\ 4. The minimum value of semi-major axis for the Oort cloud is $a_{min}$ $\\ll$ 1 $\\times$ 10$^{4}$ AU. This condition was obtained both from the numerical results on cometary evolution under the action of the galactic tides and from the observational distribution of long-period comets. If the density function of semi-major axis is approximated by proportionality $a^{\\alpha}$, then $\\alpha$ is $-$ 1/2, approximately. \\\\ 5. The magnitude of the change in perihelion distance per orbit, $\\Delta q$, of a comet due to galactic tides is a strong function of semi-major axis $a$, proportional to $a^{8.25}$. ",
        "introduction": "Intense modeling of the Oort cloud of comets started practically immediately after the Oort's paper (Oort 1950). As for the curent status of the ideas about the Oort cloud we refer to review papers, e. g., Dones et al. (2004), Levison and Jones (2007).  New physical access to the theoretical modeling of the Oort cloud of comets was suggested by Kla\\v{c}ka (2009a, 2009b). It is based on a new approach to treating the effect of galactic tides, including improvement of the model of the Galaxy. Its significance was discussed by K\\'{o}mar et al. (2009). Analytical approach to secular evolution of orbital elements was discussed by P\\'{a}stor et al. (2009). The aim of this paper is to present some characteristics obtained from the new model. The new model is represented by Eqs. (26)-(27) in Kla\\v{c}ka (2009a) and denoted as the Model II in K\\'{o}mar et al. (2009). The last paper showed the importance of the new model in comparison with the conventional models. Thus, we will not make any calculation based on the conventional models. Instead of that we will concentrate on calculations based on the new modeling of galactic tides. The results of the calculations will be compared with the results obtained by other authors isung conventional models of galactic tides.  Sec. 2 deals with distribution of cometary inclinations and it sheds some light also on the value of the exponent $\\alpha$ describing distribution of semi-major axes of comets in the Oort cloud. Sec. 3 shows that the magnitude of the change in perihelion distance per orbit of a comet due to galactic tides is a strong function of semi-major axis $a$, proportional to $a^{8.25}$. Sec. 4 relates the semi-major axis of a comet and period of oscillation in other orbital elements due to galactic tides. The found relation reminds the third Kepler's law, also in the case when natural units of length and time are used. Secs. 5 and 6 discuss some theoretical access to cometary distributions in eccentricity, perihelion distance and semi-major axis. However, Sec. 6.3 presents also results obtained by numerical integration of equation of motion for comets under the action of the Sun and the Galaxy. Sec. 7 determines minimal perihelion distances $q_{min}$ and inclinations for the instant when $q_{min}$ occurs. The section uses analytical approach to secular time derivatives of orbital elements when the initial cometary inclinations differ from 90$^{\\circ}$. The two sources of gravity, the Sun and the Galaxy, are used also in Sec. 8 dealing with the distribution of comets in inclination to ecliptic. Sec. 9 improves the most relevant previous results also for the case when the planet Jupiter gravitationally influences the motion of the comets. Finally, Sec. 10 deals with the distribution function in semi-major axis and the number of comets in the Oort cloud, including estimate of the mass of the Oort cloud. ",
        "conclusions": "The main results are: \\\\ 1. Theoretical number of long-period comets with perihelion distance $q$ $<$ 5 AU is about 50-times greater than the conventional approach yields. (Gravity of Jupiter was taken into account in finding this result.) Mass of the Oort cloud of comets is, probably, about 1/4  mass of the Earth. \\\\ 2. Semi-major axis $a$ and period of oscillations $P$ of eccentricity (and other orbital elements) are related as $a^{3}$ $P$ $=$ 1 in natural units for a moving Solar System in the Galaxy. The natural unit for time is the orbital period of the Solar System revolution around the galactic center and the natural unit for measuring the semi-major axis is its maximum value for the half-radius of the Solar System corresponding to the half-radius of the Oort cloud. The relation holds for the cases when comets approach the inner part of the Solar System, e.g., perihelion distances are less than $\\approx$ 100 AU. \\\\ 3. The minimum value of semi-major axis for the Oort cloud is $a_{min}$ $\\ll$ 1 $\\times$ 10$^{4}$ AU. This condition was obtained both from the numerical results on cometary evolution under the action of the galactic tides and from the observational distribution of long-period comets. If the density function of semi-major axis is approximated by proportionality $a^{\\alpha}$, then $\\alpha$ is, approximately, $-$ 1/2. \\\\ 4. The magnitude of the change in perihelion distance per orbit, $\\Delta q$, of a comet due to galactic tides is a strong function of semi-major axis $a$, proportional to $a^{8.25}$.  The usage of the density function of eccentricity in the form  $f(e)$ $=$ 2 $e$ cannot be argumented to be the result obtained by Jeans (1919). Consequence of the Jeans calculations is marginal density function of semi-major axis proportional to $\\sqrt{a}$ and this is not consistent with the conventionally used distributions. Moreover, Figs. 7 and 8 in K\\'{o}mar et al. (2009) show that gravitational tides alone would produce only values $e$ $<$ 0.9 for galactic inclinations $i$ less than 90 degrees and practically $e$ $<$ 0.8 for $i$ $\\approx$ 90 degrees: only for very short time intervals $e$ is greater than 0.8 ($e$ is close to 1 during the cometary visit of the inner part of the Solar System).  Using the observations yielding the distribution function of perihelion distance $F_{q} (q)$ $\\propto$ $q$ (inner part of the Solar System), we cannot come to the conclusion that  $f(e)$ $=$ 2 $e$. Any density function of the form $f(e)$ $=$ ( $k$ $+$ 1 ) $e^{k}$, $k$ $>$ 0 is also consistent with the observational result $F_{q} (q)$ $\\propto$ $q$.  Observational data on distribution in semi-major axis $a$, theoretical analytical results on the distribution functions $F_{a} (a)$ and of a perihelion distance, and, detailed numerical calculations on gravitational action of the Galaxy lead to a conclusion on $F_{a} (a)$. The conclusion is that approximation of $F_{a} (a)$ by the form proportional to $a^{\\alpha + 1}$ is consistent with all the mentioned methods if $\\alpha$ is about $-$ 1/2.  The obtained results may be improved to be more accurate. This will require also more robust and more detailed numerical calculations (and, also, several planets, not only Jupiter, have to be considered)."
    },
    "1005/1005.4691_arXiv.txt": {
        "abstract": "We study large-scale winds driven from uniformly bright self-gravitating disks radiating near the Eddington limit. We show that the ratio of the radiation pressure force to the gravitational force increases with height to a maximum of twice its value at the disk surface. Thus, uniformly bright self-gravitating disks radiating at the Eddington limit are fundamentally unstable to driving large-scale winds. These results contrast with the spherically symmetric case, where super-Eddington luminosities are required for wind formation. We apply this theory to galactic winds from starburst galaxies that approach the Eddington limit for dust. For hydrodynamically coupled gas and dust, we find that the asymptotic velocity of the wind is $v_{\\infty}\\simeq3\\langle v_{\\rm rot}\\rangle$ and that $v_{\\infty}\\propto {\\rm SFR}^{0.36}$, where $\\langle v_{\\rm rot}\\rangle$ is the mean disk rotation velocity and SFR is the star formation rate, both of which are in agreement with observations. However, these results of the model neglect the gravitational potential of the surrounding dark matter halo and a (potentially massive) old stellar bulge, which both act to decrease $v_\\infty$. A more realistic treatment shows that the flow can either be unbound, or bound, forming a ``fountain flow'' with a typical turning timescale of $t_{\\rm turn}\\sim0.1-1$ Gyr, depending on the ratio of the mass and radius of the starburst disk relative to the total mass and break (or scale) radius of the dark matter halo or bulge. We provide quantitative criteria and scaling relations for assessing whether or not a starburst of given properties can drive unbound flows via the mechanism described in this paper. Importantly, we note that because $t_{\\rm turn}$ is longer than the star formation timescale (gas mass/star formation rate) in the starbursts and ultra-luminous infrared galaxies for which our theory is most applicable, if starbursts are selected as such, they may be observed to have strong outflows along the line of sight with a maximum velocity $v_{\\rm max}$ comparable to $\\sim3\\langle v_{\\rm rot}\\rangle$, even though their winds are in fact bound on large scales. ",
        "introduction": "Galactic-scale winds are ubiquitous in starburst galaxies in both the local and high-redshift universe (\\citealt{Heckman90}; \\citealt{Heckman00}; \\citealt{Pettini01, Pettini02}; \\citealt{Shapley03}; \\citealt{rupke05}; \\citealt{Sawi08}). They are important for determining the chemical evolution of galaxies and the mass-metallicity relation (\\citealt{DS86}; \\citealt{Tremonti04}; \\citealt{Erb06}; \\citealt{finlator}; \\citealt{Peeples11}), and as a primary source of metals in the intergalactic medium (IGM; e.g., \\citealt{aguirre01}). Moreover, galactic winds are perhaps the most extreme manifestation of the feedback between star formation in a galaxy and its interstellar medium (ISM). This feedback mechanism is crucial for understanding galaxy formation and evolution over cosmic time (\\citealt{SH03}; \\citealt{OD06}; \\citealt{OD08}; \\citealt{Oppen10}).  The most well-developed  model for galactic winds from starbursts is the supernova-driven model of \\cite{CC85}, which assumes that the energy from multiple stellar winds and core-collapse supernovae in the starburst is efficiently thermalized.  The resulting hot flow drives gas out of the host, sweeping up the cool ISM (\\citealt{Heckman93}; \\citealt{SS00}; \\citealt{strickland02}; \\citealt{SH09}; \\citealt{Fujita09}). Although this model is successful in explaining the X-ray properties of starbursts, the recent observational results that galaxies with higher star formation rates (SFRs) accelerate the absorbing cold gas clouds to higher velocities ($v\\propto$SFR$^{0.35}$) and that the wind velocity is correlated with the galaxy escape velocity may challenge the traditional hypothesis that the cool gas is accelerated by the ram pressure of the hot supernova-heated wind, whose X-ray emission temperature varies little with SFR, circular velocity, and host galaxy mass, indicating a critical galaxy mass below which most of the hot wind escapes. These new observations may instead favor momentum-driven or radiation pressure-driven models for the wind physics (e.g., \\citealt{Martin99}; \\citealt{Martin05}; \\citealt{Weiner09}; Murray, Quataert \\& Thompson 2005 [hereafter MQT05]).  The model that galactic winds may be driven by momentum deposition provided by radiation pressure from the continuum absorption and scattering of starlight on dust grains was developed by MQT05. However, the conclusions of MQT05 are based on an assumed isothermal potential and spherical geometry, and are thus most appropriate for bright elliptical/spheroidal galaxies in formation. On the other hand, the theory of radiation-driven winds from accretion disks from the stellar to galactic scales has also been studied (e.g., \\citealt{TF96, TF98}; \\citealt{Proga98, Proga99}; \\citealt{Proga00, Proga03}), but none of these works considered radiation from self-gravitating disks.  In this paper we answer the question of whether or not large-scale winds can be driven by radiation pressure from self-gravitating disks radiating near the Eddington limit. These considerations are motivated by the work of Thompson, Quataert \\& Murray (2005; TQM05), who argued that radiation pressure on dust is the dominant feedback mechanism in starbursts, and that in these systems star formation is Eddington-limited. For simplicity, throughout this paper we assume that the disk is of uniform brightness and surface density. In \\S2, we show that such disks are fundamentally unstable to wind formation because the radiation pressure force dominates gravity in the vertical direction above the disk surface. This result is qualitatively different from the well-known case in spherical symmetry. In \\S2, we also discuss the applicability of this model to starbursts and then calculate the terminal velocity of the wind along the disk pole, and its dependence on both the SFR and galaxy escape velocity. In \\S3, we assess the importance of a spherical stellar bulge and dark matter halo potential. In \\S4, we discuss the 3-dimensional wind structure and estimate the total wind mass loss rate. We discuss our findings and conclude in \\S5. ",
        "conclusions": "In this paper we study the large-scale winds from uniformly bright self-gravitating disks radiating near the Eddington limit. Different from the spherical case, where the Eddington ratio $\\Gamma=f_{\\rm rad}/f_{\\rm grav}$ is a constant with distance from the source, for disks $\\Gamma$ increases to a maximum of twice its value at the disk surface (\\S2). As a result, such disks radiating at (or even somewhat below; see Fig.~\\ref{fig_subEdd}) the Eddington limit are unstable to driving large-scale winds by radiation pressure.  We quantify the characteristics of the resulting outflow in the context of Eddington-limited starbursts, motivated by the work of TQM05 who argued that radiation pressure on dust is the dominant feedback process in starburst galaxies. We find that the asymptotic terminal velocity along the polar direction from disks without a stellar bulge or dark matter halo is $v_{\\infty}\\sim\\sqrt{4\\pi G\\Sigma r_{D}}\\sim3\\langle v_{\\rm rot}\\rangle$, where $r_{D}$ is the disk radius, $\\Sigma$ is the surface density, $v_{\\rm rot}$ is the disk rotation velocity (see eqs.~\\ref{terminal01} and \\ref{velocity_1}), and may range from $\\sim50-1000$\\,km s$^{-1}$ for starbursts, depending on the system considered. Furthermore, by employing the observed Schmidt law, we find that $v_{\\infty}\\propto$ SFR$^{0.36}$r$_{D}^{-0.21}$ (see eq.~\\ref{terminal02}). These results, in the absence of dark matter halo and bulge, are in agreement with recent observations (e.g., Martin 2005, 2006; Weiner et al.~2009; see also Chen et al.~2010). The typical mass loss rate from an Eddington-limited disk in the single-scattering limit is given by equation (\\ref{massloss02}) and suggests these outflows may efficiently remove mass from the disk (eq.~\\ref{massloss03}).  However, both wind velocities and outflow rates can be significantly decreased by the presence of a spherical old stellar bulge or dark matter halo potential (\\S3). Deeper or more extended spherical potentials cause the flow to be bound on large scales, and to produce only ``fountain flows'' where particles fall back to the disk on a typical timescale of $\\sim0.1-1$ Gyr, depending on the parameters of the system considered (see Figs.~\\ref{fig_halo}-\\ref{fig_velocity_1}). The criterion for the flow to become bound is given in equation (\\ref{f0limt}). For typical values of the parameter $f_0$ (see eq.~\\ref{scalef0}), we find that the winds from starbursts can be either bound or unbound (see Fig.~\\ref{fig_halo}). For $f_{0}\\simeq0.25$ in equation (\\ref{scalef0}), the asymptotic velocity of the wind is decreased by a factor of $\\simeq0.7$ from the case neglecting the dark matter halo completely (compare eqs.~\\ref{velocity_1} and \\ref{decreas02}), from $\\simeq3\\langle v_{\\rm rot}\\rangle$ to $\\simeq2.3\\langle v_{\\rm rot}\\rangle$.  An increase in $f_0$ by a factor of 2 would decrease the asymptotic velocity further to $\\simeq1.3\\langle v_{\\rm rot}\\rangle$, whereas an increase of $f_0$ by a factor of $\\sim3$ to $\\sim 0.75$ would cause the flow to become bound.  Importantly, even in the limit of bound fountain flows, if the timescale $t_{\\rm turn}$ for reaching the turning point $z_{\\rm turn}$ is longer than the lifetime of the starburst $t_{\\star}$ (eq.~\\ref{tstar}), one may still observe an outward going wind while the starburst is active and bright. The maximum positive velocity of the flow $v_{\\rm max}$ along the line of sight for non-edge-on disks is still correlated with $v_{\\rm rot}$ of the disk (see \\S3). These facts may complicate the inference from observations of winds that they are unbound from the surrounding large-scale dark matter halo.  Clearly, more detailed work is required to fully assess radiation pressure on dust as the mechanism for launching cool gas from starburst disks. The main limitations of the work presented here are that (1) we consider only disks with constant brightness and surface density, (2) we do not compute the hydrodynamics of the flow, but instead treat the outflow in the test-particle limit, and (3) we ignore other physical effects such as the momentum and energy input from supernovae (Chevalier \\& Clegg 1985; Strickland \\& Heckman 2009), which may act in concert with radiation pressure (MQT05; Murray et al.~2010). The general case, with hydrodynamics, realistic disk brightness, surface density, and dust opacity profiles will modify the picture presented here. Such an effort is underway (Zhang \\& Thompson, in preparation). However, the basic conclusion that uniformly bright self-gravitating disks radiating near the Eddington limit are able to drive large-scale winds --- particularly in the high-$\\Sigma$ limit in starbursts (see \\S2)--- should not be fundamentally changed by more elaborate considerations. Indeed, although we have specialized the discussion to starbursts and dust opacity, the instability derived in \\S2 is of general applicability."
    },
    "1005/1005.4658_arXiv.txt": {
        "abstract": "The altitude distribution of optical turbulence is derived from the MASS instrument data by solving an inverse problem. In this paper, some modifications of the profile restoration are described. The principal change is the introduction of the Non Negative Least Squares algorithm which has good regularizing properties. An averaging of scintillation indices was replaced with averaging of obtained solutions what leads to clearer physical results. It is shown that restoration with a number of turbulent layers as large as 14--15  can be successfully performed. ",
        "introduction": "It is well known that the efficiency of astronomical observations in optical and near IR range greatly depends on turbulence in the earth's atmosphere. Modern techniques for telescope efficiency gain require both statistically-valid long-term and near real-time information about properties of the optical turbulence (OT) above an observatory. In general, the turbulence intensity is described with help of the refraction index structure constant $C_n^2$. One of the instruments designed to measure this parameter is MASS (Multi Aperture Scintillation Sensor) \\cite{Marr} developed a decade ago. More than 30 such devices are used in different astronomical observatories and site testing campaigns  \\cite{Kor07}.  The method used by the MASS \\cite{Marr,MASS} to determine the OT vertical profile is based on the simultaneous measurements of stellar scintillations in four entrance apertures ($A$, $B$, $C$, $D$) of different size and produces four normal and six differential \\cite{T98} scintillation indices $s^2$ --- a variance of the relative fluctuations of light fluxes. The measurement is performed with an integration time $T_e$ of 1~ms and an estimation of $s^2$ is calculated from the measurements at time interval $T_b$ ({\\it basetime}) of about 1~s \\cite{MASS,mnras2003}.  Scintillation indices $s^2$ are affected by turbulence strength in the whole atmosphere traveled through by the light: \\begin{equation} s^2 = \\int_0^\\infty C_n^2(h) Q(h)\\,dh, \\label{eq:sc} \\end{equation} where $C_n^2(h)$ is index constant at altitude $h$, $Q(h)$ is certain weighting function that depends on geometry of light propagation: size of receiving aperture and angular size of emitting object. The function $Q(h)$ can be calculated for a given geometry and stated spatial spectrum of refraction index perturbation \\cite{MASS,mnras2003}.  The calculation of the vertical distribution of the OT from measured $s^2$ indices is only possible by solving the inverse problem which is ill-conditioned for any practical set of MASS apertures (therefore, set of $Q(h)$ functions). For this reason, the restoration algorithm of the turbulence profiles is a key part of the MASS method. The description of the algorithm and its verification are given in paper \\cite{mnras2003}. Instrumental and other factors affecting accuracy of MASS results were considered in \\cite{Kor07,mnras2007}.  During development of the algorithm used in the {\\it Turbina} software (hereafter --- algorithm P), some simplifications were adopted with intuitive guesses  due to lack of real data for careful analysis of the algorithm work.  In recent years, intensive measurements with the MASS instrument have been performed by different groups at different sites. The large amount of data encouraged us to revise the profiles restoration technique. The revision focuses on the implementation of a more valid mathematical standpoint and on the verification (and correction if it is necessary) of the model assumptions.  A particularly significant factor for us is the completion of a three year campaign at mount Maidanak \\cite{maid2005} and a successful two year campaign at Shatdzhatmaz summit \\cite{mnras2010} where 2.5~m telescope of Sternberg institute should be installed. ",
        "conclusions": "The described modifications of the OT profiles restoration process do not, in principle, change the physical results obtained with MASS instrument. However, the revised restoration algorithm makes the results more reliable and clearer for interpretation. Studies of the profile restoration were carried out before in connection with MASS instrument verification \\cite{mnras2003,mnras2007,apo}. Many discovered effects stimulated the work presented in this paper.  We formulate a summary of the modifications made to the algorithm and their consequences on the results obtained with the MASS technique.  \\begin{enumerate}  \\item Algorithm based on direct minimization is replaced by a more mathematically correct method of Non-Negative Least Squares (NNLS).  \\item Comparative analysis of the restoration of altitude profiles of the optical turbulence on measurements at mount Maidanak in 2005 -- 2007 shows that the NNLS algorithm works always better.  \\item Calculation of solution errors has shown that in the case of weak turbulence errors are comparable with the turbulence intensities.  \\item The consequence of that is the distortion of the statistical distribution of layer intensities, leading to a significant underestimation of the median and quartile values.  \\item The main source of uncertainty is the non-stationarity of OT at the scale of the integration time (usually 60~s), noise of the measurements is almost always much lower.  \\item De-correlation of the input data errors slightly changes the estimates of output errors and does not change the solutions themselves.  \\item The change of the strategy of the restoration algorithm from ``averaging indices -- restoration''  to ``restoration -- averaging profiles'' has led to a significant decrease of the errors of the OT profiles.  \\item As a result, vertical profiles have more statistical significance and physical clarity. The integrated turbulence intensity in free atmosphere has increased by no more than 15 \\%.  \\item In the case of turbulence, corresponding to good seeing, it is possible to increase the vertical resolution and accuracy of turbulent layer localization (from $ h / 2 $ to $ h / 4 $) using the  12 -- 14 layers model.  \\item The new algorithm is also effective for data obtained with the MASS/DIMM instrument, which has  different set of the entrance apertures.  \\end{enumerate}  An additional consequence from the revision of restoration process is the improvement of the C++ code that is clearer and more suitable for further modifications and a substantial (order of magnitude) increase in processing speed."
    },
    "1005/1005.3287_arXiv.txt": {
        "abstract": "A simple set of diagrammatic rules is formulated for perturbative evaluation of ``in-in\" correlators, as is needed in cosmology and other nonequilibrium problems.  These rules are both intuitive, and efficient for calculational purposes. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.1885_arXiv.txt": {
        "abstract": "{Classical T~Tauri stars (CTTS) are young, late-type objects, that still accrete matter from a circumstellar disk. Analytical treatments and numerical simulations predict instabilities of the accretion shock on the stellar surface.}{We search for variability on timescales below a few minutes in the CTTS TW~Hya and AA~Tau.}{TW~Hya was observed with SALTICAM on the Southern African Large Telescope (SALT) in narrow-band filters around the Balmer jump. The observations were performed in slit mode, which provides a time resolution of about 0.1~s. For AA~Tau we obtained observations with OPTIMA, a single photon-counting device with even better time resolution.}{Small-scale variability typically lasts a few seconds, however, no significant periodicity is detected. We place a 99\\% confidence upper limit on the pulsed fraction of the lightcurves. The relative amplitude is below 0.001 for TW~Hya in the frequency range 0.02-3~Hz in the 340~nm filter and 0.1-3~Hz in the 380~nm filter. The corresponding value for AA~Tau is an amplitude of 0.005 for 0.02-50~Hz.}{ The relevant timescales indicate that shock instabilites should not be seen directly in our optical and UV observations, but the predicted oscialltions would induce observable variations in the reddening. We discuss how the magnetic field could stabilise the accretion shock.} ",
        "introduction": "\\label{introduction} Classical T~Tauri stars (CTTS) are young ($<10$~Myr), low-mass ($<3M_{\\sun}$) pre-main sequence objects. They still actively accrete from a circumstellar disk. The disk does not reach down to the central star, but is truncated by the stellar magnetic field approximately at the co-rotation radius \\citep{1983ards.proc..625U,1988ApJ...330..350B,1991ApJ...370L..39K}. The high-energy photons from the star irradiate, heat and ionise the inner parts of the disk, while the matter is funneled along the magnetic field lines and impacts on the star at close to free-fall velocity. For the simple case of a dipolar field, the accretion spots are located at high stellar latitude \\citep{1991ApJ...370L..39K,1994ApJ...429..781S}. A strong shock develops on or near the stellar surface, which heats up the accreted matter to a few $10^6$ K. Observational support for this model comes from different wavelength bands. The disk is seen in molecular hydrogen emission lines \\citep{2002ApJ...572..310H}, and the inner holes can be resolved interferometrically with radio observations \\citep{2006ApJ...637L.133E}. The accretion funnel emits strongly in H$\\alpha$ \\citep{1998ApJ...492..743M}. This is one of the defining characteristics of CTTS. The accretion spot on the stellar surface is hot enough to emit UV and X-ray radiation \\citep{lamzin, acc_model}. It also heats the underlying stellar photosphere to about $20\\;000$~K \\citep{calvetgullbring}.  This variety of different phenomena leads to time variability on different scales. Long changes occur over years, possibly related to a slowly changing circumstellar extinction \\citep{2007A&A...461..183G}. For many CTTS the dominant source of optical variability is the stellar rotation with a typical period of a few days \\citep{1983ApJ...267..191R,1986A&A...165..110B,2009ApJ...694L.153N}. Depending on its inner radius, the rotation timescale of the inner disk and also the free-fall time from the disk onto the star is a few hours. Thus variability in the accretion rate should be limited to this timescale; indeed, variations of about 0.1~mag are regularly observed in broad-band photometry from CTTS \\citep{1996MNRAS.282..167S,2008MNRAS.391.1913R}.  Nevertheless, much faster variations can be expected if the accretion shock on the stellar surface is unstable. This has first been studied in the context of accretion onto white dwarfs \\citep{1983ASSL..101..199L,1982ApJ...261..543C,1985ApJ...296..128I}. Even if the magnetic field is strong enough to channel the accretion stream and prevent turbulence on the boundary between accretion stream and surrounding photosphere, the shock can still fluctuate up and down. The infalling matter builds up a funnel above the stellar surface. Owing to the ram pressure of the infalling matter, the gas gets compressed and accumulates. At some point the accumulated matter is dense enough for a runaway cooling to set in and the column collapses fast. A new shock forms at the stellar surface and works its way upwards again building up a new reservoir of matter. \\citet{2008A&A...491L..17S} and \\citet{2009arXiv0912.1799O} predict a typical timescale of 10~min for the fluctuations, and the same is seen in models of \\citet{2008MNRAS.388..357K}, but with much shorter periods of 0.02-0.2~s. The difference can be mainly attributed to different cooling functions and different input densities $n$. Since the cooling is proportional to $n^2$, different pre-shock densities heavily influence the cooling time. \\citet{2008MNRAS.388..357K} predict the shock luminosity to fluctuate about 50\\% around the mean, but caution that this value and the exact shape of the fluctuation are very model dependent. They also find that magnetic fields tend to stabilise the accretion shock, so that no oscillations occur.  The most direct tracer of the accretion shock is the soft X-ray emission. This is formed just behind the shock front, where the accreted material is heated to a 2-3 MK \\citep{lamzin,acc_model}. \\citet{2009ApJ...703.1224D} used \\emph{XMM-Newton} X-ray data to search for periodicity in \\object{TW Hya}, one of the best-studied CTTS. They did not detect any period and set an upper limit on the pulsed fraction of 5\\% at the 99\\% confidence level over the interval 0.001-7~Hz. TW~Hya is the only target where the X-ray count rate is high enough to perform timing analyses on scales below a minute. In this paper we use a secondary tracer of the accretion shock. We obtained observations in the blue band, using optical telescopes. In the blue and the UV range, CTTS are dominated by the continuum emission from the accretion shock \\citep{1974MNRAS.168..603L,1980Afz....16..243G,2000ApJ...535L..47M} because they are late-type objects and their photospheric emission quickly drops towards shorter wavelengths. It is not clear from the current models how fast the optical emission reacts to shock instabilities, but observations in the visible potentially provide a diagnostic with high signal-to-noise ratio, and yield a complementary view on sub-minute variability.  We will first introduce the targets observed in Sect.~\\ref{targets} and the observations performed in Sect.~\\ref{observations}. Details of the data reduction process are given in Sect.~\\ref{datareduction}. In Sect.~\\ref{periodicity} we search for periodicities in the data employing Lomb-Scargle periodograms. We discuss our results in Sect.~\\ref{discussion} and summarise them in Sect.~\\ref{summary}. ",
        "conclusions": "\\label{discussion} Our analysis of the data presented above yields no significant periodicity in the searched frequency range. The limits set by this data on the maximum amplitude for variations are up to two orders of magnitude more sensitive than those \\citet{2009ApJ...703.1224D} derived for the X-ray lightcurve of TW~Hya. However, the soft X-rays are formed directly in the shock cooling zone with a weak contribution from the stellar corona \\citep{acc_model}, therefore much stronger oscillations are expected in X-rays \\citep{2008A&A...491L..17S} than in the UV wavelength range we probe. So far no time-dependent radiative-transfer calculations of accretion shocks in CTTS have been carried out. In stationary 1-D models the hot gas at the lower end of the accretion column and the heated photosphere below  emit a significant portion of their energy around the Balmer-jump \\citep{calvetgullbring}. In this wavelength region the accretion spot dominates  over the stellar photosphere by one order of magnitude in TW~Hya \\citep{2000ApJ...535L..47M}. The heated photosphere is optically thick and can be treated as a black-body with temperatures between 5000~K and 8000~K \\citep[see Fig.~3 in][]{calvetgullbring}, its luminosity should thus be of the order of $10^{11}$~erg~s$^{-1}$~cm$^{-2}$.  The \\citet{calvetgullbring} model refers to younger CTTS with a radius of $2 R_{\\sun}$, however the spectral type of TW~Hya matches the model, thus we can apply it to obtain an order-of-magnitude estimate. The model implies a column mass of the heated photosphere down to the region, where the continuum is formed, of the order 10~g~cm$^{-2}$, containing a thermal energy of $10^{13}$~erg~s$^{-1}$~cm$^{-2}$, and suggests a timescale for cooling of the heated photosphere of 100~s. This calculation gives only a rough estimate, because the cooling would change the temperature and thus the opacity structure, so the deeper layers become visible as the cooling proceeds. It suggests that rapid changes of the shock conditions will not be visible in optical and UV emission from the heated photosphere; however, detailed radiative transfer calculations are required to confirm this.  According to the models the flow time of accreted matter through the post-shock zone is a few seconds. Incidentally, the flow time through the shock cooling zone matches the time we find in the autocorrelation function for TW~Hya. The short term variability in the lightcurves shown in figures~\\ref{TW_Hya_lc_34}-\\ref{AATau_lc} is dominated by statistical fluctuations, so the source variability, which is detected in the autocorrelation function, is of the order of 1\\% of the continuum level at most. Several effects could be responsible for that: If the accretion flow is inhomogeneous and clumpy, this would cause the UV emission from the shock cooling zone to vary within seconds. The size of the accreted clumps is presumably small in comparison to the total area covered by accretion spots, which is estimated from 0.3\\% to 5-10\\% of the stellar surface for TW~Hya, depending on the wavelength of the observation \\citep[][and references therein]{acc_model}. If the clumps switch on or off the emission from 1\\% of the accretion spot, then their characteristic size is $10^3-10^4$~km. The sound speed in the hot parts of the post-shock cooling zone is of the order 100~km~s$^{-1}$, thus an area with radius $10^3$~km could be affected within a few seconds; in cooler regions of the post-shock cooling zone, where most of the optical emission originates, the temperature and thus the sound speed are about an order of magnitude lower. In this region only very small areas can respond coherently to disturbances.  Instead of inhomogeneous accretion the sound waves could also be caused by accretion shock instabilities. Three-dimensional simulations are needed to study this in more detail. But would the optical emission react fast enough to changing conditions, however they are caused? Even if it does not, the oscillations of the accretion shock could be seen in absorption. In the time-dependent simulations of \\citet{2008A&A...491L..17S} the accretion shock oscillates up and down. Thus, the pre-shock column density changes with time by $\\approx10^{20}$~cm$^{-2}$. This accreted material is not reprocessed in the hot shock yet and we apply the standard conversion between X-ray column density and optical absorption $A_V= N_H/ 2\\times 10^{21}$~cm$^{-2}=0.05$~mag \\citep{2003A&A...408..581V}. According to \\citet{1999PASP..111...63F} an additional absorption of $A_V=0.05$~mag would reduce the flux at 3800~\\AA{} by 7\\%, much more than our upper limit on periodic variability. The numbers given are calculated for a line-of-sight along the accretion flow. TW~Hya is seen pole-on, so this condition is valid for accretion at high latitude. The extinction changes with the viewing geometry; the argument does not apply if the accretion column is seen from the side.  While general variability is seen, clear periods are absent with significant limits down to 0.001 for the maximum amplitude of a pulsed fraction of the emission. It might simply be hard to observe oscillations in the optical, as discussed above, but \\citet{2009ApJ...703.1224D} also do not see a periodicity in the X-ray lightcurve of TW~Hya. Those authors present a detailed discussion on the expected period of the oscillations, the effect of multiple accretion streams and the influence of the magnetic field, arguing the supporting pressure of the magnetic field has to be taken into account.  Different studies of time-dependent accretion shocks concentrate on different aspects. \\citet{2008MNRAS.388..357K} present 1D magneto-hydrodynamic models of super-Alfv\\'enic accretion flows with infall velocities around 100~km~s$^{-1}$ and different incident angles. In their simulations strong magnetic fields lead to a stable accretion flow. They develop a stability criterion which depends on the angle between the field and the stellar surface. However, this criterion has not been verified for higher infall velocities. If it is still valid for velocities of 525~km~s$^{-1}$ and densities of $10^{-12}$~g~cm$^{-3}$, as measured from X-ray spectroscopy in TW~Hya \\citep{2002ApJ...567..434K,twhya,acc_model,2009A&A...505..755R,2010ApJ...710.1835B}, then a magnetic field of 100~G is sufficient to keep the shocks stable for an accretion flow perpendicular to the stellar surface.  Yet, it is not clear if this is also valid for the stronger magnetic fields usually measured in the accretion spot regions of CTTS \\citep{2004Ap&SS.292..619V, 2008MNRAS.386.1234D},  which indicate sub-Alfv\\'enic flows. \\citet{2010A&A...510A..71O} performed 2D simulations of accretion perpendicular to the stellar surface including heat conduction  in both the super-Alfv\\'enic and the sub-Alfv\\'enic regime. They confirm the importance of the magnetic field, and find that shocks with high $\\beta$ (the ratio between gas pressure and magnetic pressure) become stable through magnetic damping  for super-Alfv\\'enic flows. \\citet{2010A&A...510A..71O} observe quasi-periodic shock oscillations for cases with larger magnetic fields  when the flow becomes sub-Alfv\\'enic and chaotic behaviour without periodicity for even weaker magnetic fields. \\citet{2010A&A...510A..71O} also compare their results to 1D simulations with a similar setup by \\citet{2008A&A...491L..17S} and find that 2D models predict a smaller amplitude and higher frequency than 1D models. For large values of $\\beta$ the magnetic field cannot confine the accretion stream to a funnel and matter flows out side-ways, disturbing the surrounding stellar atmosphere.  Although \\citet{2008MNRAS.388..357K} and \\citet{2010A&A...510A..71O} predict stability in different regimes, they both agree that there are magnetic field configurations that inhibit shock oscillations in the super-Alfv\\'enic regime. Only \\citet{2010A&A...510A..71O} probe sub-Alfv\\'enic flows, but they test less parameters than \\citet{2008MNRAS.388..357K}, so we cannot exclude that similar stability regions exit for sub-Alfv\\'enic flows. So,  we may not see periodic oscillations of the accretion shock, not because of confusion of a larger number of accretion streams, but because the shock is intrinsically stable."
    },
    "1005/1005.5624_arXiv.txt": {
        "abstract": "{We study the upward propagation of a localized velocity pulse that is initially launched below the transition region within the solar atmosphere. The pulse quickly steepens into a shock, which may lead to the formation of spicules. } {We aim to explore the spicule formation scenario in the framework of the rebound shock model. } {We solve two-dimensional time-dependent magnetohydrodynamic equations numerically to find spatial and temporal dynamics of spicules. } {The numerical simulations show that the strong initial pulse may lead to the quasi periodic rising of chromospheric material into the lower corona in the form of spicules. The periodicity results from the nonlinear wake that is formed behind the pulse in the stratified atmosphere. The superposition of raising and falling off plasma portions resembles the time sequence of single and double (sometimes even triple) spicules, which is consistent with observational findings. } {The two-dimensional rebound shock model may explain the observed speed, width, and heights of type I spicules, as well as observed multi-structural and bi-directional flows. The model also predicts the appearance of spicules with $3-5$ min period due to the consecutive shocks.  } ",
        "introduction": "Although spicules were discovered almost 130 years ago (Secchi \\cite{secchi1877}), they still remain one of mysterious phenomena in the solar atmosphere. Spicules are usually detected in chromospheric H$_\\alpha$, D$_3$, and Ca II H lines as thin and elongated structures in the solar limb. Many authors consider H$_\\alpha$ dark mottles as a disc counterpart of limb spicules (Suematsu et al. \\cite{Suematsu1995}), however this point is still under discussion. Excellent reviews of the general properties of spicules (and mottles) were presented almost forty years ago by Beckers (\\cite{bec68,bec72}). Suematsu (\\cite{Suematsu1998}) and Sterling (\\cite{Sterling2000}) published more recent reviews of observational and theoretical aspects of spicules, respectively, while Zaqarashvili \\& Erd{\\'e}lyi (\\cite{Zaqarashvili2009}) summarized observed oscillation events. Very recently, Pasachoff et al. (\\cite{pas09}) studied high-resolution dynamics of limb spicules by Swedish Solar Telescope (SST) with 0.2 arc sec resolution. High resolution observations by Solar Optical Telescope (SOT) on board of the recently launched Hinode show another type of spicules, which have quite different properties than classical limb spicules (De Pontieu et al. \\cite{dep07b}). The classical and newly observed spicules are referred to as types I and II, respectively.  Disc counterparts of limb spicules still are not known despite of all the discussion over the years (Tsiropoula et al. \\cite{Tsiropoula1994}, Sterling \\cite{Sterling2000}). However, recent high-resolution observations on SST and Hinode/SOT suggest that active-region dynamic fibrils (Hansteen et al. \\cite{hansten2006}, De Pontieu et al. \\cite{dep07}) and quiet Sun mottles (Rouppe van der Voort et al. \\cite{rouppe2007}) have properties similar to type I spicules. On the other hand, Rouppe van der Voort et al. (\\cite{rouppe2009}) report discovering of disc counterparts to type II spicules through observations on SST.  Despite all the mechanisms of spicule formation proposed in the literature none of them provides convincing explanations of all properties. Spicule formation mechanisms can be formally divided into three different groups: pulses, Alfv\\'en waves, and p-mode leakage.  The main idea of spicule formation by impulsively launched perturbations is as follows. Velocity or gas pressure pulses deposited in the lower atmosphere are steepened into shocks as a result of the rapid decrease in mass density with height. These shocks lift up the transition region, producing spicules. Hollweg (\\cite{hol82}) developed the shock model with one-dimensional simulations. In his rebound shock model, Hollweg simulated gas pressure to lift the transition region to the observed altitudes. Suematsu et al. (\\cite{Suematsu1982}) did the same for velocity pulses. Later on, Sterling et al. (\\cite{Sterling1990,Sterling1993}), Cheng (\\cite{cheng1992}), and Heggland et al. (\\cite{heg07}) studied the shock models, including radiation and heat conduction.  Hollweg et al. (\\cite{holw82}) showed that Alfv\\'en waves may be nonlinearly coupled to fast magnetoacoustic shocks, which may lift up the transition region. Later on, Kudoh \\& Shibata (\\cite{kud99}) suggested that the random nonlinear Alfv\\'enic pulses may reproduce the spicules at the observed heights. On the other hand, the damping of high-frequency Alfv\\'en waves due to ion-neutral collisions may be responsible for formation of spicules (Haerendel \\cite{Haerendel1992}, James et al. \\cite{james03}).  De Pontieu et al. (2004) propose that, as a result of fall off of acoustic cut-off frequency, p-modes may be channeled into the solar corona along inclined magnetic field lines. The oscillations then may be steepened into shocks producing spicules. De Pontieu et al. (2004) argue that the observed quasi $5$-min period in spicule appearance is associated with the periodicity of p-modes. Several papers (Hansteen et al. \\cite{hansten2006}, De Pontieu et al. \\cite{dep07}) then studied the formation of spicules due to the leakage of photospheric convective motions and oscillations into upper atmosphere. By performing numerical simulations of MHD equations the authors show that oscillations and flows generated by self consistent two-dimensional (2D) convective motions drive spicules.  Numerous other mechanisms of spicule formation were suggested, such as resonant buffeting of anchored magnetic flux tubes by solar granules (Roberts \\cite{roberts1979}). Detailed summary of these mechanisms can be found in the review paper by Sterling (\\cite{Sterling2000}). Horizontal magnetic flux emergence has also been found to be one of many possible spicule drivers, the others being convective overshoot, collapse of granules, p-modes, and reconnection in either the photosphere or chromosphere (Mart\\'inez-Sykora et al. \\cite{martinez09}).  Despite significant achievements in the above-mentioned papers, there are still no existing mechanisms that explain the observed double structures of spicules (Tanaka \\cite{Tanaka1974}, Dara et al. \\cite{Dara1998}, Suematsu et al. \\cite{Suematsu2008}) and bi-directional flow (Tsiropoula et al. \\cite{Tsiropoula1994}, Tziotziou et al. \\cite{Tziotziou2003,Tziotziou2004}, Pasachoff et al. \\cite{pas09}). Suematsu et al. (\\cite{Suematsu2008}) and Tsiropoula et al. (\\cite{Tsiropoula1994}) have proposed a magnetic reconnection as an explanation of, respectively, the double structures and bi-directional flow. However, neither cogent numerical simulations nor analytical models have been presented so far. Therefore, the problem remains still open to discussion.  In this paper we perform 2D numerical simulations of magnetohydrodynamic (MHD) equations and show that the 2D rebound shock model may explain both the double structures and bi-directional flow in spicules. Unlike previous 2D numerical simulations (Hansteen et al. \\cite{hansten2006}, De Pontieu et al. \\cite{dep07}), which considered a periodic initial driver, we launch a single initial pulse and follow its temporal development. Such an approach allows us to trace propagation of the pulse and the internal structure of a spicule in detail, which explains double structure and bi-directional flow.  This paper is organized as follows. The main properties of limb spicules are summarized in Sect.~2. A numerical model is presented in Sect.~3, and the corresponding numerical results are shown in Sect.~4. This paper concludes with a discussion and short summary of the main results in Sects.~5 \\& 6, respectively. ",
        "conclusions": "We performed 2D numerical simulations of velocity pulse propagation in the solar atmosphere. The strong longitudinal pulse was launched at the lower part of the chromosphere, and its consequent propagation along transition region into the corona was traced. The pulse quickly steepens into shock, which influences the dynamics of the upper chromosphere/transition region; namely, the rarefied tail behind the shock front leads to the pressure gradient above the transition region, which lifts up the chromospheric material in the form of spicule. The strong initial pulse may raise the chromospheric plasma up to observed heights ($6-7$ Mm) and with observed speed ($25$ km s$^{-1}$). The spicule begins to falls back after some time, while the secondary shock (rebound shock) lifts up another portion of the chromospheric material. The superposition of falling off and rising plasma portions produces the observed double structure and bi-directional flows in spicules. The simulated spicule exhibits a temporal development with double and single structures (sometime even triple) that are akin to the observational data. The model predicts quasi-periodic raising of spicules with the period of the nonlinear wake that is formed behind a leading pulse in the stratified atmosphere. The periodicity strongly depends on the initial amplitude of the pulse and can be in the range of $3-5$ min and even longer. We believe that future sophisticated numerical simulations may give a more complete picture of spicule formation in the frame of the 2D and 3D rebound shock models."
    },
    "1005/1005.4764_arXiv.txt": {
        "abstract": "{Searches for exoplanets with radial velocity techniques are increasingly sensitive to stellar activity. It is therefore crucial to characterize how this activity influences radial velocity measurements in their study of the detectability of planets in these conditions. } {In a previous work we simulated the impact of spots and plages on the radial velocity of the Sun. Our objective is to compare this simulation with the observed radial velocity of the Sun for the same period. } {We use Dopplergrams and magnetograms obtained by MDI/SOHO over one solar cycle to reconstruct the solar integrated radial velocity in the Ni line 6768 \\AA. We also characterize the relation between the velocity and the local magnetic field to interpret our results. } {We obtain a stronger redshift in places where the local magnetic field is larger (and as a consequence for larger magnetic structures): hence we find a higher attenuation of the convective blueshift in plages than in the network. Our results are compatible with an attenuation of this blueshift by about 50\\% when averaged over plages and network. We obtain an integrated radial velocity with an amplitude over the solar cycle of about 8 m/s, with small-scale variations similar to the results of the simulation, once they are scaled to the Ni line. } {The observed solar integrated radial velocity agrees with the result of the simulation made in  our previous work within 30\\%, which validates this simulation. The observed amplitude confirms that the impact of the convective blueshift attenuation in magnetic regions will be critical to detect Earth-mass planets in the habitable zone around solar-like stars. } ",
        "introduction": "Searches for exoplanets with radial velocity (hereafter RV) techniques will be increasingly sensitive to stellar activity as the sensitivity of instruments is improved \\cite[such as Espresso on the VLT or Codex on the E-ELT][]{dodorico07}. It is therefore crucial to understand how this activity influences the RV measurements. In our previous works, we simulated the influence of spots \\cite[][hereafter Paper I]{lagrange09} and plages \\cite[][hereafter Paper II]{meunier09} on the radial velocity (hereafter RV) for the Sun seen equator-on. The simulation covered a solar cycle with an excellent temporal cadence. Our study considered the photometric contribution of spots and plages and also the attenuation of the convective blueshift in magnetic regions (Paper II). This last effect appeared to be dominant, with a long-term amplitude over the solar cycle of about 8~m/s, while the photometric contributions produced a rms of 0.08 m/s during cycle minimum and of 0.42 m/s during cycle maximum (0.33~m/s over the whole cycle), with peak-to-peak amplitudes of a few m/s. We studied how these RV variations influence the detectability of an Earth-mass planet in the habitable zone and found that it made this detection very difficult.  Most existing observations of the solar RV do not cover our data set. Neither are the investigated lines well characterized in terms of convective blueshift, which makes them difficult to compare with simulations. \\cite{mcmillan93} measured the solar RV with the moon light observed in the violet part of the spectrum for the period 1987--1992. They found solar RV variations that are significantly smaller than our simulation. For the period 1983--1992, \\cite{DP94} measured the solar RV in the 2.3~$\\mu$m domain while \\cite{jimenez86} measured it in the potassium line at 769.9~nm \\cite[][]{brookes78} during cycle 21 (1976--1984). Both have found an amplitude over the cycle of about 30~m/s. The potassium line is a strong line however, for which a particularly high convective blueshift is not expected. The GOLF \\cite[][]{gabriel95} instrument on SOHO provides integrated measurements of the solar RV in the Na D1 and D2 lines. Unfortunately, the planned calibration procedure could not be applied to the GOLF observations because of an instrumental failure early in the mission. Therefore the long-term variations of the RV signal and its amplitudes are unreliable (P. Boumier, private communication). We note that after a few years of operation, a comparison has been made between the MDI and GOLF signal \\cite[][]{henney98} but focused on the high temporal frequency component and not on the long-term signal. Some attempts to calibrate the long term variations a posteriori, mostly due to the spacecraft orbit, have been made \\cite[][]{ulrich00} but the resulting signal still shows some strong instrumental residuals that are not compatible with the amplitude we wish to test.   This is why precise RV observations of the Sun in integrated light are necessary for testing our simulations, in particular for the convective blueshift contribution. It would be best to directly observe the Sun in integrated light over an extended domain of the spectrum with an instrument and data analysis similar to those used to measure stellar RV variations. This is a long-term project though. The aim of this paper is to characterize the solar RV variations with spatially resolved observations by MDI/SOHO in a single spectral line (Ni line at 6768 \\AA). The use of an imaging instrument also allows us to identify the contribution of solar activity to the solar integrated velocities, in contrast with an instrument such as GOLF. We reconstruct the integrated RV variations (neglecting the photometric component of spots and plages simulated in Paper I and II) and characterize the behavior of the velocity variations measured in the Ni line with respect to the magnetic field. Our observations are compared to the simulations performed in Paper II scaled to the expected behavior of that line.  The input data and data analysis are described in Sect.~2. We describe how we determine the zero velocity for each data point in the time series and estimate the uncertainty of our measurements. The local velocity - magnetic field relation derived from the analysis of the Ni line is presented in Sect. 3. This relation allows the interpretation of the integrated RV variations. We also discuss the origin of the velocity associated to magnetic regions in term of convective blueshift attenuation and downflows. Then we present the integrated RV in Sect.~4 and compare it to the results of the simulation of the convective blueshift attenuation made in Paper II scaled to the Ni line expected behavior. The comparison is made at different time-scales. We conclude in Sect.~5. ",
        "conclusions": "We reconstructed the integrated RVs of the Sun from MDI Dopplergrams using the Ni line at 6768 \\AA$\\:$ over a solar cycle. We obtained a long-term amplitude of the variations of about 8~m/s. The resulting RV variations agree very well with the simulated RV (Paper II) once scaled to the Ni line: the small-scale and long term variations are well reproduced within 30\\%. We therefore validate the simulations made in Paper II.  We also analyzed in detail the variations of the relative velocity (with respect to the quiet Sun) in magnetic regions, which we interpreted as an attenuation of the convective blueshift in most cases. We characterized this velocity as a function of the local magnetic field and as a function of the structure size. Our results are compatible with an average attenuation of the convective blueshift in plages and network regions by about 50 \\%, compared to a factor 2/3 in \\cite{bs90} for plages. The difference probably arises because we included weaker field regions in our analysis.  The variations of the RV due to the convective blueshift attenuation versus the structure size could be implemented in the simulations like that made in Paper II, because the structure size is one of the input parameters defining magnetic regions. The dependence on the local magnetic field is important to understand and test our RV variations. However, if we want to extrapolate this work to other stars, this relation may not be easy to take into account, because the magnetic field distribution on other stars is not well characterized."
    },
    "1005/1005.1188_arXiv.txt": {
        "abstract": "\\vspace{.2cm} \\noindent Slow-roll inflation is studied in theories where the inflaton field is conformally coupled to the Ricci scalar. In particular, the case of Higgs field inflation in the context of the noncommutative spectral action is analyzed. It is shown that while the Higgs potential can lead to the slow-roll conditions being satisfied once the running of the self-coupling at two-loops is included, the constraints imposed from the CMB data make the predictions of such a scenario incompatible with the measured value of the top quark mass. We also analyze the r\\^ole of an additional conformally coupled massless scalar field, which arises naturally in the context of noncommutative geometry, for inflationary scenarios. ",
        "introduction": "Cosmological inflation is the most widely accepted mechanism to resolve the shortcomings of the standard Hot Big Bang model. This mechanism, leading to a phase of exponential expansion in the very early universe, is deeply rooted in the fundamental principles of General Relativity and Field Theory, and once combined with the principles of Quantum Mechanics, it can account for the origin of the observed large scale structures and the measured temperature anisotropies of the Cosmic Microwave Background (CMB). However, despite its success, cosmological inflation remains a paradigm in search of a model which should be motivated by a fundamental theory. The strength of the inflationary mechanism is based on the assumption that its onset is generically independent of the initial conditions. Nevertheless, even this issue is under debate~\\cite{Calzetta:1992gv,Calzetta:1992bp,Trodden:1999wc,Gibbons:2006pa,Germani:2007rt,Ashtekar:2009mm} given the lack of a complete theory of Quantum Gravity.  The inflaton field (usually a scalar field) is assumed to dominate the evolution of the universe at early times, but its origin and the form of its effective potential both remain unknown; for this reason it would be attractive if the one scalar field that {\\sl is} commonly thought to exist, namely the Higgs field, also doubled as the long-searched for inflaton.  Unfortunately, it seems that if the Higgs field is minimally coupled to gravity this cannot be achieved, which has led some authors to consider large nonminimal couplings of the Higgs field to gravity where inflation might be achieved~\\cite{Bezrukov:2007ep}.  It is commonly assumed/chosen that there is no coupling ({\\sl i.e.}, minimal coupling) between the inflaton field and the background geometry (the Ricci curvature). However, this assumption/choice seems to lack a solid justification. A first (and merely aesthetic) motivation comes from the observation that in the early universe (where masses are negligible), the equations of motion for spinors and gauge bosons have a natural conformal invariance in four space-time dimensions, while the same is true for scalar fields only when they couple to the Ricci scalar in a specific way.  More compelling is the fact that even if classically the coupling between the scalar field and the Ricci curvature could be set equal to zero, a nonminimal coupling will be induced once quantum corrections in the classical field theory are considered.  Moreover, a nonminimal coupling seems to be needed in order to renormalize the scalar field theory in a curved space-time.  The precise value of the coupling constant (denoted by $\\xi$) then depends on the choice of the theory of gravity and the scalar field~\\cite{Faraoni:1996rf}.  It has also been argued that in all metric theories of gravity, including General Relativity, in which the scalar field is not part of the gravitational sector ({\\sl e.g.}, when the scalar field is the Higgs field), the coupling constant should be conformal in order for the short distance propagators of the theory to match those found in a Minkowski space-time --- a requirement of the strong equivalence principle~\\cite{Faraoni:1996rf,Sonego:1993fw} (in our notation, conformal coupling means $\\xi=1/12$). Finally, in the context of finite theories at one-loop level, it was shown~\\cite{Buchbinder:1989bt,Yoon:1996yr} that the nonminimal coupling $\\xi$ tends either to its conformal value or increases exponentially in modulus, depending on the specific structure of the theory.  In what follows, we will investigate whether scalar fields, and in particular the Higgs field, could play the r\\^ole of the inflaton in the presence of a small positive nonminimal coupling between the scalar field and the background geometry. The coupling constant $\\xi$ is not a free parameter which could be tuned to achieve a successful inflationary scenario avoiding severe fine-tuning of inflationary parameters ({\\sl e.g.}, the self-coupling of the inflaton field), $\\xi$ should instead be dictated by the underlying theory. For negative values of $\\xi$, exponential expansion is more easily achieved than in the minimal case, and it can in fact lead to inflation consistent with observational data in the strong coupling limit~\\cite{Bezrukov:2007ep,DeSimone:2008ei}.  In fact, the slow-roll parameters for large $|\\xi|$ are independent of $\\xi$ and only depend on the number of e-folds. However, exponential expansion is less favoured for positive values (in our conventions) such as conformal coupling~\\cite{Faraoni:2000wk}.  In light of the motivations for a small positive $\\xi$ outlined above, we will investigate whether quantum corrections to the Higgs potential can lead to a slow-roll inflationary era and if so, whether the constraints imposed from the CMB temperature anisotropies are satisfied.  We will apply this analysis to the Spectral Action of NonCommutative Geometry (NCG). This theory leads naturally to a Lagrangian with a conformal coupling between the Higgs field and the background geometry, in the form of a boundary condition at high energies $E\\geq\\Lambda$, where $\\Lambda$ is a characteristic scale of the model. NCG provides an elegant way of accounting for the Standard Model (SM) of Particle Physics and its phenomenology~\\cite{Chamseddine:2006ep}. Our motivation is to investigate cosmological consequences of the NCG Spectral Action and, in particular, to test whether slow-roll inflation driven by one of the scalar fields arising naturally within NCG could be realized in agreement with experimental data and astrophysical measurements.  In a previous study, we (one of us and a collaborator) have studied~\\cite{Nelson:2009wr} the conditions on the couplings so that the Higgs field could play the r\\^ole of the inflaton in the context of the NCG.  Since however the running of couplings with the cut-off scale had been only analyzed~\\cite{Chamseddine:2006ep} neglecting the nonminimal coupling between the Higgs field and the curvature, we were not able to reach a definite conclusion. In this respect, the study below is a follow-up of Ref.~\\cite{Nelson:2009wr}. Moreover, it has been argued~\\cite{Marcolli:2009in} that inflation with a conformally coupled Higgs boson could be realized in the context of NCG due to the running of the effective gravitational constant. In what follows, we will also analyze the validity of this statement. Finally, the NCG Spectral Action provides, in addition to the Higgs field, another conformally coupled (massless) scalar field, which exhibits no coupling to the matter sector~\\cite{Chamseddine:2009yf}.  One may {\\sl a priori} wish/expect that this field could be another candidate for the inflaton; we will examine this scenario as well.  Concluding, we analyze slow-roll inflation within models that exhibit a conformal coupling between the Higgs field and the Ricci curvature. Our motivation is to investigate whether any of the two scalar fields arising naturally within the NCG Spectral Action could be identified as the inflaton. As we will explicitly show, our analysis leads us to the conclusion that unfortunately such a slow-roll inflationary scenario fails to remain in agreement with current data from high energy physics experiments and astrophysical measurements.  This paper is organized as follows: In Section~\\ref{slow-roll}, we study the issue of the realization of slow-roll inflation within theories with a nonminimal coupling between the scalar field and the Ricci curvature, classically. The analysis is first performed in the Jordan frame and then in the Einstein frame. In Section~\\ref{quantum}, we consider corrections to the Higgs potential through two-loop renormalization group analysis of the minimally coupled Standard Model; we then enlarge this study in the case of a conformal coupling. We focus on the gravitational and Higgs field sector of the Lagrangian density, obtained within the noncommutative spectral action, which has a conformal coupling, in Section~\\ref{NCG}. We find that even though we can accommodate an era of slow-roll inflation, it seems difficult to reach an agreement with the CMB data. This conclusion holds not only for the Higgs field but also for the other scalar field which appears generically in the theory. We then examine in Section~\\ref{running-gravitational}, whether running of the gravitational constant could modify our conclusions with regards to the realization of a successful inflationary scenario driven through one of the scalar fields in the NCG theory. We round up our conclusions in Section~\\ref{conclusions}.  Our signature convention is $(-\\,+\\,+\\,+)$; the Riemann and Ricci tensors are defined as \\beq R^\\sigma{}_{\\mu\\nu\\rho}&=&\\Gamma^\\sigma_{\\mu\\rho,\\nu}-\\Gamma^\\sigma_{\\nu\\rho,\\mu} +\\Gamma^\\tau_{\\mu\\rho}\\Gamma^\\sigma_{\\tau\\nu}- \\Gamma^\\tau_{\\nu\\rho}\\Gamma^\\sigma_{\\tau\\mu}~,\\nonumber\\\\ R_{\\mu\\nu}&=&R^\\rho{}_{\\mu\\rho\\nu}~,\\nonumber \\eeq respectively. Note that within our definition of $\\xi$, conformal coupling means $\\xi=1/12$. ",
        "conclusions": "\\label{conclusions} In many realistic cosmological models, the nonminimal coupling of the scalar field to the Ricci curvature cannot be avoided. In particular, there are arguments requiring a conformal coupling between the scalar field and the background curvature. The existence of such a term will generically lead to difficulties in achieving a slow-roll inflationary era. In this paper, we have investigated whether two-loop corrections to the Higgs potential could lead to a slow-roll inflationary period in agreement with the constraints imposed by the CMB measurements. Our findings do not favor the realization of such an era. More precisely, even though slow-roll inflation can be realized, we cannot satisfy the COBE normalization constraint for any values of the top quark and Higgs masses allowed from current experimental data.  We have in particular investigated Higgs inflation in the context of the Noncommutative Geometry Spectral Action, which provides an elegant explanation for the phenomenology of the Standard Model. Within this context, a conformal coupling arises naturally between the Higgs field and the Ricci curvature.  It is also important to note that once conformal coupling is set at the preferential (boundary) energy scale of the spectral action model, then it will remain conformal at all scales.  Running of the gravitational constant and corrections by considering the more appropriate de\\,Sitter, instead of a Minkowski, background do not favor the realization of a successful inflationary era.  The NCG Spectral Action provides in addition to the Higgs field, another (massless) scalar field which exhibits no coupling to the matter sector. Our analysis has shown that neither this field can lead to a {\\sl successful} slow-roll inflationary era if the coupling values are conformal. One may be able to improve upon this (negative) conclusion, if important deviations of $\\xi$ from its conformal value can be allowed; the value $\\xi=1/12$ may turn out that is not a generic feature of NCG models.  \\bigskip"
    },
    "1005/1005.3054_arXiv.txt": {
        "abstract": "{ We describe the Herschel Virgo Cluster Survey (HeViCS) and the first data obtained as part of the Science Demonstration Phase (SDP). The data cover a central 4$\\times$4 sq deg region of the cluster. We use SPIRE and PACS photometry data to produce 100, 160, 250, 350 and 500 $\\mu$m luminosityfFunctions (LFs) for optically bright galaxies that are selected at 500 $\\mu$m and detected in all bands. We compare these LFs with those previously derived using IRAS, BLAST and Herschel-ATLAS data. The Virgo Cluster LFs do not have the large numbers of faint galaxies or examples of very luminous galaxies seen previously in surveys covering less dense environments.} ",
        "introduction": "The Virgo Cluster provides us with a unique opportunity to study in detail a large number of galaxies in the cluster environment. Virgo is probably the most studied cluster of galaxies because of its proximity to the Milky Way - it lies at a distance of $\\sim$17 Mpc (Gavazzi et al., 1999 and references therein) with a mean velocity of $\\sim$1094 km s$^{-1}$ (Binggeli et al., 1987). It is an Abell richness Class I cluster containing $\\sim$2000 optically catalogued galaxies (Virgo Cluster Catalogue (VCC), Binggeli et al., 1985). The Virgo Cluster virial mass is $\\approx 2.5 \\times 10^{14}$ M$_{\\odot}$ (Girardi et al., 1998; Rines \\& Diaferio 2006) as expected for a cluster as rich as this. The cluster contains large amounts of X-ray emitting gas (Boehringer et al., 1994) and shows clear evidence for both sub-structure and non-virialised motions (Gavazzi et al., 1999). From previous observations of galaxy clusters we know that this environment can have a significant effect on the properties of galaxies. The morphological mix is quite different compared to the field and many galaxies may be deficient in atomic gas (Haynes \\& Giovanelli, 1984). This deficiency is probably responsible for the lower star-formation rates (Lewis et al., 2002, Gomez et al., 2003), the truncation of star-forming discs (Boselli \\& Gavazzi, 2006) and the observed higher metal content of cluster galaxies (Skillman et al., 1996). Being so nearby, the Virgo Cluster has previously been studied in detail at many wavelengths and so there is a wealth of data available, to which we are now able to add the FIR.  This is the first of a series of seven papers, based on HeViCS SDP data, in which we discuss: how the cluster environment truncates the dust discs of spiral galaxies (paper II, Cortese et al., 2010), the dust life-time in early-type galaxies (paper III, Clemens et al., 2010), the spiral galaxy dust surface density and temperature distribution (paper IV, Smith et al., 2010), the properties of metal-poor star-forming dwarf galaxies (paper V, Grossi et al., 2010), the lack of thermal emission from the elliptical galaxy M87 (paper VI, Baes et al., 2010) and the far-infrared (FIR) detection of dwarf elliptical galaxies (paper VII, De Looze et al., 2010). A further paper discusses the spectral energy distributions of HeViCS galaxies together with  others observed as part of the Herschel Reference Survey (Boselli et al., 2010). In this paper we briefly discuss the HeViCS project, describe the data used in the above papers and compare the Virgo FIR LFs with those previously derived.  \\begin{figure} \\centering \\includegraphics[scale=0.45]{vcc_hevics.epsi} \\caption{ Virgo Cluster. The solid black line indicates the area covered by the VCC, with each VCC galaxy marked by a small black circle. The ROSAT X-ray contours are marked in blue. The area to be observed by HeViCS is indicated by the four red boxes. The SDP data described here are from the field centred on M87.} \\end{figure}  \\vspace{-0.2cm} ",
        "conclusions": "We have described the HeViCS and our first science data from PACS and SPIRE. The data were used to produce 100, 160, 250, 350 and 500 $\\mu$m LFs for a sample of optically bright galaxies in the central region of the Virgo Cluster. Unlike the global optical LF and those previously derived in the FIR these functions do not appear to have a steep power law dependence at the faint end. There are relatively few faint FIR sources that can be associated with confirmed Virgo Cluster members compared to what is found for field galaxies and there are no examples of very luminous FIR sources.   \\\\"
    },
    "1005/1005.1098_arXiv.txt": {
        "abstract": "Feedback from protostellar outflows can influence the nature of turbulence in star forming regions even if they are not the primary source of velocity dispersion for all scales of molecular clouds. For  the rate and power expected in star forming regions, we previously \\citep{Carroll2009} demonstrated that outflows could drive supersonic turbulence at levels consistent with the scaling relations from \\cite{matzner2007} although with a steeper velocity power spectrum than expected for an isotropically driven supersonic turbulent cascade. Here we perform higher resolution simulations and combine simulations of outflow driven turbulence with those  of isotropically forced turbulence. We find that the presence of outflows within an ambient isotropically driven turbulent environment produces a knee in the  velocity power spectrum at the outflow scale and a steeper slope at sub-outflow  scales than for a purely isotropically forced case. We also find that the presence of outflows flattens the density spectrum at large scales effectively reducing the formation of large scale turbulent density structures. These effects are qualitatively independent of resolution.  We have also carried out Principal Component Analysis (PCA) for synthetic data from our simulations. We find that PCA as a tool for identifying the driving scale of turbulence has a misleading bias toward low amplitude large scale velocity structures even when they are not necessarily the dominant energy containing scales. This bias is absent for isotropically forced turbulence but manifests strongly for collimated outflow driven turbulence. ",
        "introduction": "\\paragraph{} Supersonic velocity dispersions are a ubiquitous phenomena occuring over a wide range of scales in star forming molecular clouds.  The relationship between velocity dispersions (line-widths) and size was first summarized by \\citep{larson}.  While Larson's laws originally applied only to global relationships between different GMC's, \\cite{heyer} showed that the same line-width size relationships held within individual GMC's.  These observed velocity dispersions are commonly interpreted as resulting from turbulent gas motion.  The formal distinction between turbulent motion and other sources of velocity dispersion in molecular clouds such as gravitational fragmentation is an important one \\citep{field,fbk10},  but for simplicity we presently use the terms turbulence and velocity dispersions interchangeably. \\paragraph{} The source of turbulent motions in molecular clouds is a topic of debate.  If self-gravity is ignored then large scale sources of energy for turbulence include the kinetic energy injection from supernova and galactic differential rotation. Sources of turbulent energy on smaller scales include stellar feedback in the form of HII regions, radiatively driven winds or accretion driven outflows.   This raises the issue of the so-called injection scale, and whether there are in fact more than one dynamically important injection scale where energy is deposited into the turbulent spectrum.  For a clump scale cloud ($\\le  10$pc),  theoretical scenarios tend to divide the sources of turbulent forcing mentioned above into external and internal:  External forcing models consider the injection scale to be larger than the cloud scale (e.g. Supernova driven ISM turbulence) whereas internal forcing would refer to injection scales smaller than the cloud size (HII regions, outflows)\\citep{McKeeOstriker} or winds from B stars \\citep{Valverde2010}.  Note that if self-gravity is included, then gravitational contraction from very large scales can itself produce velocity dispersions that are consistent with Larson's laws \\citep{field,fbk10}. That represents an alternative paradigm for driving turbulent-like motions in molecular clouds. Whatever the source of the motions, in order for the turbulence in the parsec and subparsec scale clumps and cores to remain in an approximate steady state, it needs to be continually driven over time scales longer than a crossing time either internally via stellar feedback or externally via a cascade down from the ISM. \\paragraph{} Distinct from the question of the origin of turbulent motions, is the role of turbulent 'feedback' on the star formation process.  In the absence of turbulent support, most of the mass within a given structure (GMC, clump, core) would collapse into stars within one free fall time giving an efficiency per free fall time $\\epsilon_{ff}=1$ \\citep{McKeeOstriker}.  Observations however yield surprising low values of $\\epsilon_{ff}$ ranging from .01 to .1 not only for GMC's but also for substructures within GMC's of mean densities $n_H$ ranging from $10^2$ to $10^5$ cm$^{-3}$ \\citep{krumholz-turb-reg,krumholztan}.  One explanation for these low values of $\\epsilon_{ff}$ is the that the observed turbulent motions are signatures of internal kinetic support against collapse. We emphasize that the global energy budget of turbulence is a separate issue from the the issue of local turbulent feedback. For example, star clusters form on scales that are a fraction of the total molecular cloud scale $L_{cloud}/L_{cluster} \\sim 10$ and dynamical feedback from outflows can impede star formation on cluster forming scales even if the total energy budget from outflows is insufficient to power the observed turbulent energy budget of the entire GMC.  This means that internal sources of turbulent forcing may be important on small scales whether or not they are important on large scales.  On even smaller scales (ie. $L_{core} \\sim .1$ pc) outflows are likely responsible for the dispersal of gas \\citep{VelusamyLanger1998, MatznerMckee2000, Myers2008} especially if the outflow opening angle widens with time \\citep{ArceSargent2006}, although the complex interaction between cores and their outflows requires more numerical modeling.  In addition, not all forcings are equal: some forcings are isotropic and some are not. This distinction has observational consequences. \\paragraph{} In this paper we focus on the physics and observational consequences of collimated outflow driving into both initially laminar and initially turbulent ambient environments.  Theoretical treatments of feedback from multiple spherical outflows creating a self-regulating star forming system was first explored by \\cite{normansilk}.  Analytical work by Matzner \\citep{matzner00,matzner01,matzner2007} has explored the role of collimated outflow feedback on clouds. \\cite{krumholz} also considered the nature of feedback via outflows, concluding that these systems provide an important source of internal driving in dense star forming cores. The first simulation based study of the problem was presented by \\cite{maclow-outflows}.  More recently \\cite{linaka} and \\cite{nakamura} have mapped out the complex interplay between star formation outflows and turbulence.  Their simulations include a self-consistent formulation of driving outflows from newly formed stars and they concluded that protostellar outflows were a viable means of generating turbulence in star forming clusters. \\cite{Wang2010} built on these studies using higher resolution AMR studies demonstrating the vigorous feedback that outflows are able to provide a cluster including the significant modification of the star formation efficiency and development of massive stars.  We note that the Nakamura, Li and Wang studies all focused on cluster scales rather than molecular cloud scales. In addition they did not focus specifically on the mechanisms by which outflows generated turbulence or the physics of power spectra signatures.  We also note a study of single jets by \\cite{banerjee} that explored the volume filling averages of post shock concluding that single jets would not leave enough supersonic material in their wakes to act as a relevant source of internal forcing. \\paragraph{} The study of the detailed mechanics of outflow driven turbulence was taken up in a series of papers by \\cite{cunningham-jc,cunningham-cavity,cunningham-sj}. These investigations built off observational work by \\cite{quillen} for NGC 1333 which showed that turbulence was likely driven not by active outflows but by the interaction of fossil shells driven by extinct outflows.  In \\cite{cunningham-sj} it was shown how single outflows in a turbulent environment would be disrupted leading to their energy resupplying turbulent motions. In \\cite{Carroll2009} a study of multiple, randomly oriented jets driven into a quiescent medium was undertaken. The results of these simulations showed conclusively that interacting outflows can lead to turbulent motions.  Explorations of the energy spectrum and the density PDF showed that statistically steady isotropic random motions were maintained by the outflows.  Of particular interest was the fact that an injection scale associated with outflow cavity interaction $K_{outflow}$ was apparent in the energy spectrum $E(k)$.  Also the power law index of the spectrum was steeper than would be expected for Burgers type turbulence such as that generated when less realistic isotropic forcing at large scales is used. \\cite{Carroll2009} found that the steeper spectrum was associated with outflow cavities sweeping up vortices at smaller scales.  {\\it Thus the Carroll results provided both proof that outflows can drive turbulence and provided some insights into the physics involved and its statistical signatures}. \\paragraph{} The goal of discerning scales of  turbulent driving in real molecular clouds has, recently, received considerable attention.  In \\cite{bruntheyermaclow} PCA was used to determine the dominant scales of turbulent driving based on molecular line emission.  Based on comparisons of  first and second eigenimages of molecular channel maps for $^{13}$CO and $^{12}$CO they find that the data support the conclusion that that the bulk of the energy driving for turbulence comes from scales much larger than that of clumps.  \\cite{padoan-vca} reached a similar conclusion by applying VCS (Velocity Coordinate Spectrum) \\citep{lazarian09} on $^{13}$CO data.   \\cite{bruntheyermaclow} also analyzed C$^{18}$0 data which probes denser regions of the clumps, and this data did reveal evidence for a second, subparsec  driving scale consistent with that of outflows.  While  neither of these studies preclude the importance of feedback from outflows on very small scales, neither of these papers actually invoke full 3-D simulations of collimated outflow driven turbulence in their conclusions about driving scale.  Thus there is  a potential loophole in the application of methods that require a density-velocity scaling relation in order to interpret observations as a signature of kinetic energy density on a given scale. One of the topics we address in this paper is the bias that results from this specific issue. \\paragraph{} In section 2 we discuss the basic parameters and setup of our simulations and choices of forcing for both our collimated outflow driving and our isotropically forced cases.  In section 3 we discuss the results comparing the isotropically driven turbulence with collimated outflow driven turbulence and the interplay between the two. In particular, we analyze the growth and saturation of turbulence, the velocity and density spectra, the characteristic driving scales, and the relative insensitivity to resolution. In section 4 we uncover biases in the current use of PCA  to infer the scales of turbulent driving when comparing isotropically forced to outflow driven turbulence.  We show that PCA fails to identify where the dominant driving power is when comparing sources of turbulence with different forcings.  We conclude in section 5. ",
        "conclusions": "\\label{conclusion} Building on the work of \\cite{Carroll2009} we studied the properties of turbulence driven isotropically, by outflows, and by a combination of both for input parameters typical of clump scale molecular clouds.  We find as in \\cite{Carroll2009} that outflows are able to drive supersonic turbulence at levels consistent with the scaling relations put forth by \\cite{matzner2007}. We also find that turbulence driven by outflows is characterized by a steeper velocity spectrum at sub-outflow scales, and an overall flatter density spectrum.  Based on our resolution study, we conclude that when outflows are present, the resulting steep velocity spectra and correspondingly smooth velocity structures seen at sub-outflow scales are not due to numerical dissipation, but rather due to the presence of the outflows themselves.  Outflows are able to continually generate smooth velocity structures on outflow scales that sweep up and 'dissipate' smaller structures at the same rate that these structures would begin to cascade.  Outflow driven turbulence is therefore quite different from the cascade of energy produced by large scale isotropic forcing. \\paragraph{} Since outflows drive motions on sub-parsec scales, it seems likely that realistic outflow driving will take place in the presence of a externally driven cascade, due to external energy inputs into molecular cloud systems from much larger scales.  With this motivation, we drove outflows into a medium also supplied with isotropcially driven turbulence from larger scales.  We found that the properties of the turbulence at sub-outflow scales were largely unchanged by the external isotropic cascade for typical cloud parameters. In particular, the velocity spectra still contain a knee at the outflow scale and a steep slope at sub-outflow scales.  We also found that the outflows produced a flatter density spectra even at supra-outflow scales.  We attribute this to the disruption by outflows of larger scale structures produced by the large scale driving.  We conclude therefore, as in \\cite{Wang2010}, that outflows have the potential to play a large role in setting the star formation rate and the initial mass function in their neighborhoods. \\paragraph{} Finally, we also find that the density-velocity distributions and to a smaller extent correlations present in outflow driven turbulence are quite different from those present in an isotropic cascade.  Many of the techniques for inferring the velocity spectrum of turbulence in star forming regions rely on assumptions about density-velocity distributions and correlations based on simulations of isotropically forced turbulence.  To illustrate the resulting potential bias, we  tested one such model (PCA) against the results from our simulations of outflow driven turbulence.  We find that indeed PCA underestimates the contribution of outflows to the velocity dispersion, though not necessarily the mean scalar momentum.  We find that PCA is able to infer the driving scale of the outflows for the case of pure outflow driven turbulence, but that the presence of even small velocity large scale flows quickly leads to an overestimation.  We suggest weighting the data cubes by the line of sight velocity as one way to make PCA more sensitive to higher velocity flows.  This highlights that nature of the forcing of turbulence (e.g. isotropic on large scales vs. outflow driven) is of direct importance for interpreting observations."
    },
    "1005/1005.1051_arXiv.txt": {
        "abstract": "\\small{\\\\ GRB afterglow 090510 is (so far) the best-monitored afterglow in the optical, X-ray, and above 100 MeV, measurements covering 2--3 decades in time at each frequency. Owing to its power-law temporal decay and power-law spectrum, it seems very likely that the highest energy emission is from the forward-shock energizing the ambient medium (the standard blast-wave model for GRB afterglows), the GeV flux and its decay rate being consistent with that model's expectations. However, the synchrotron emission from a collimated outflow (the standard jet model) has difficulties in accounting for the lower-energy afterglow emission, where a simultaneous break occurs at 2 ks in the optical and X-ray light-curves, but with the optical flux decay (before and after the break) being much slower than in the X-rays (at same time). The measured X-ray and GeV fluxes are incompatible with the higher-energy afterglow emission being from same spectral component as the lower-energy afterglow emission, which suggests a synchrotron self-Compton model for this afterglow. Cessation of energy injection in the blast-wave and an ambient medium with a wind-like $n \\propto r^{-2}$ density can explain all features of the optical and X-ray light-curves of GRB afterglow 090510. Such an ambient medium radial structure is incompatible with this short-GRB originating from the merger of two compact stars. } ",
        "introduction": "The existence of the afterglow emission following Gamma-Ray Bursts (GRBs) has been predicted by Paczy\\'nski \\& Rhoads (1993) and M\\'esz\\'aros \\& Rees (1997a) and has been been monitored in the radio (catalog of Frail et al 2003), optical (catalog of Kann et al 2010, Swift-UVOT catalog of Oates et al 2009, Roming et al 2009), and X-ray (catalogs: BeppoSAX -- De Pasquale et al 2006, Swift -- O'Brien et al 2006, Willingale et al 2007) for hundreds of afterglows and, more recently, it has been detected above 100 MeV in a few cases.  The general afterglow picture is that the relativistic outflow that produced the prompt (burst) emission interacts with the ambient medium, driving two shocks: one into the circumburst medium (the \"forward\" shock), leading to a progressive deceleration of the blast-wave) and one into the GRB ejecta (the \"reverse\" shock). The shocks accelerate electrons to relativistic energies (up to at least 10 GeV, in the comoving frame) through the first-order Fermi process or by the plasma two-stream instability (Medvedev \\& Loeb 1999). Magnetic dissipation in a Poynting outflow has also been proposed as the origin of relativistic electrons in GRB outflows (M\\'esz\\'aros \\& Rees 1997b, Lyutikov 2006). Turbulence in the shocked fluid or the mentioned Weibel instability are believed to be the origin of the $\\sim 1$ Gauss magnetic field required to explain the X-ray afterglow emission with synchrotron emission from the forward-shock.  The radio, optical, and X-ray afterglow emission has been identified most often with synchrotron emission from the ambient medium energized by the forward-shock (e.g. M\\'esz\\'aros \\& Rees 1997a, Sari, Piran \\& Narayan 1997). The reasons for that is that, for an impulsive GRB ejecta release, most of the ejecta energy is quickly transferred to the forward-shock, and that the forward-shock emission flux is expected to decay as a power-law for long times (as observed in GRB afterglows) without any further assumptions/requirements: just the power-law deceleration of the blast-wave and the power-law spectrum of the shock-accelerated electrons suffice to yield a synchrotron flux with a power-law decay in time.  The emission from the GRB ejecta energized by the reverse-shock has also been proposed to be the afterglow source, more likely at optical frequencies (M\\'esz\\'aros \\& Rees 1997a, Panaitescu \\& M\\'esz\\'aros 1998), and was identified as the origin of the bright optical counterpart of GRB 990123 (Sari \\& Piran 1999). Due to the fast cooling of the ejecta electrons, the reverse-shock emission may account for the afterglow emission lasting for days only if new electrons are accelerated at that shock, i.e. only if there is a continuous influx of ejecta crossing the reverse-shock. In this case, the afterglow light-curve also depends on the rate at which the new ejecta cross the reverse-shock, a fact used by Uhm \\& Beloborodov (2007) and by Genet, Daigne \\& Mochkovitch (2007) to explain the plateaus displayed by the X-ray afterglow light-curves at 1-10 ks. As shown in those articles, the reverse-shock emission may also account for the various degree of coupling between the optical and X-ray afterglow light-curves, although the reason for chromatic X-ray light-curve breaks lies in the tracking of the reverse-shock electron distribution and not in some fundamental/simple afterglow process.  In this work, we model the broadband emission of afterglow 090510 by calculating the synchrotron emission from both shocks. Kumar \\& Barniol Duran (2010) and Corsi et al (2010) have identified the emission of this afterglow at all wavelengths with synchrotron from the forward-shock. Analytically and numerically, we find that, if that were true, then the X-ray emission after the 2 ks break would be too dim for the GeV flux measured by Fermi-LAT at 100 s. To solve this incompatibility, the optical and GeV emissions must be from different emission processes, and we propose that the latter is inverse-Compton scatterings. The fast decay of the X-ray flux after the 2 ks break indicates that the X-ray is also inverse-Compton. As shown by Panaitescu \\& Kumar (2000), inverse-Compton scatterings in the forward-shock could be the dominant emission process in the X-ray if the ambient medium is denser than about 100 protons/${\\rm cm^{-3}}$ but, so far, there is only one case (000926 -- Harrison et al 2001, Panaitescu \\& Kumar 2002) where the X-ray afterglow light-curve and flux warranted its explanation with the synchrotron self-Compton model. ",
        "conclusions": "\\hspace*{5mm}{\\sl Uniform jet}. The salient properties of afterglow 090510 -- a slow rise of the optical flux, an {\\sl achromatic} light-curve break at 2 ks, followed by a flux decay that is substantially slower in the optical than at X-rays -- cannot be explained with the traditional jet model (synchrotron emission from the forward-shock driven by a collimated outflow into the ambient medium). The smaller problem with this model is that the optical flux rise requires the synchrotron spectral peak to be above optical before the 2 ks light-curve break, which is incompatible with the soft optical spectrum measured by UVOT at 1 ks. The bigger problem is the inability of the jet model to yield different decays at different frequencies, as measured for the optical and X-ray light-curves of afterglow 090510 after the 2 ks break (when power-law decay indices differ by $\\alpha_{x2} - \\alpha_{o2} \\simeq 1$).  {\\sl Structured outflow}. In order for the emission from different parts of the outflow to be dominant at only one frequency (above the peak energy of the synchrotron spectrum), the slope of the synchrotron spectrum must not be the same. Therefore, the optical and X-ray flux decays could be decoupled if the forward-shock accelerates electrons to a power-law distribution with energy whose index $p$ is a function of e.g. the shock's local Lorentz factor. A potential difficulty for this model is that the decoupling of the light-curve decays at two different frequencies is not sustainable for a long time, with one part of the outflow becoming eventually dominant at both frequencies, due to the continuous decrease of the synchrotron peak frequency. In the particular case of afterglow 090510, a \"top-hat\" jet model does not work because the electron indices required for each outflow region (by the optical and X-ray flux decays which they should produce) are not sufficiently different.  {\\sl Synchrotron self-Compton}. Some X-ray afterglows display a light-curve plateau at 1--10 ks whose end is chromatic and not due to a spectral break crossing the X-rays, which rules out a common origin of the optical and X-ray afterglow emissions in at least some afterglows. This may suggest that the X-ray light-curve plateau of afterglow 090510 may not be the forward-shock emission. However, even if only the X-ray measurements after the plateau are attributed to the forward-shock, a successful synchrotron forward-shock emission model cannot be found for afterglow 090510. The problem is that the X-ray flux extrapolated back in time to 100 s and in frequency to GeV is smaller by a factor of a few than the GeV flux measured at 100 s. This shows that the lower and higher frequency emissions of GRB afterglows 090510 cannot arise from the same spectral component, which prompted us to consider the synchrotron self-Compton (SsC) forward-shock model, where the lower energy (optical) afterglow emission is synchrotron and the higher frequency (GeV) emission is up-scattering of the synchrotron photons.  The only remaining ingredient is a reason for the achromatic light-curve break seen at 2 ks. Because achromatic breaks must arise from the dynamics of the outflow and because a jet-break would yield a synchrotron optical flux decay $F_{o2} \\propto t^{-p}$ that is too steep, we have considered a wide outflow for which the achromatic 2 ks break arises from cessation of energy injection in the forward-shock.  Analytically, it is expected that the inverse-Compton flux from an adiabatic forward shock decays approximately as $F_x \\propto t^{-(p-1.3)}$ for a homogeneous ambient medium and as $F_x \\propto t^{-p}$ ($t^{-(p-1)}$) for a wind medium, at frequencies below (above) the cooling frequency of the inverse-Compton spectrum. Therefore, if the ambient medium were homogeneous, the expected flux decay would be too slow compared to the post-break, fast X-ray light-curve decay of afterglow 090510. For a homogeneous medium, the SsC model also has difficulties in accounting for the relative intensity of the optical and X-ray fluxes. However, the multiband fluxes of afterglow 090510, as well as their measured decays, are well accommodated by a wind-like medium. The best-fit obtained with this model is not without flaws (GeV flux decay is slightly too slow, post-break optical flux decay is slightly too fast) and some of its features are a bit extreme (jet energy is higher by a factor 10 than for other afterglows, wind is denser by a factor 10 than the average measured for Galactic Wolf-Rayet stars).  For all the numerical calculations whose results were presented above, we find that the forward-shock emission is brighter than the reverse-shock's at all frequencies of interest (optical, X-ray, GeV) if the microphysical parameters for electron and magnetic field energy are the {\\sl same} behind both shocks. (This conclusion applies mostly to the cases where there is energy injection into the blast-wave, i.e. when there is a reverse-shock during the afterglow phase). Modelling the multiwavelength emission of afterglow 090510 with only the reverse-shock emission has failed to yield acceptable fits, either because the synchrotron model emission is well below the measured GeV flux or because it exceeds the measured optical flux when the GeV data are fit with inverse-Compton emission. However, the reverse-shock emission depends on parameters quantifying the history of ejecta release, thus the parameter space to be investigated is larger than for the forward-shock, thus we do not rule out a possible origin of the GRB afterglow 090510 in the reverse-shock energizing some incoming ejecta.  {\\sl GRB 090510: a short-burst from a massive star}. If the SsC + energy-injection is the correct explanation for afterglow 090510, then the radial stratification of the ambient medium is incompatible with this short burst originating from the merger of two compact stars. Furthermore, for the best-fit outflow parameters (kinetic energy per solid angle and ambient density), the lack of a jet-break in the synchrotron optical light-curve requires a jet energy above $10^{52}$ erg, which may exceed what the merger of two compact stars can deposit into an ultrarelativistic outflow.  Instead, the required type of medium and outflow energetics indicate that GRB 090510 is a short burst arising from a massive star. If so, it must have occurred within the host galaxy, yet the center of that host (at $z=0.9$, identified by Rau et al 2009) is at 9 kpc (in projection) from the afterglow position (McBreen et al 2010). That is very puzzling, yet not unprecedented: GRB 070125 is a long-burst (hence, it is likely to have originated also from the core-collapse of a massive star) which, probably, occurred in a young star-forming cluster located far ($> 27$ kpc) from the host center and well into its halo (Cenko et al 2008). Alternatively, GRB 090510 occurred within its host galaxy but that host is too dim to be detected (although its redshift cannot be higher than $z=1.5$, given the detection of the afterglow in near-UV bands by UVOT -- Kuin et al 2009).  A short burst arising from a massive stellar progenitor is not the traditional view about the origin of bursts shorter than 2 s. The usual dichotomy -- short GRBs are from binary mergers \\& long GRBs from core-collapse of massive stars -- is based partly on the correlation of burst timescale with those expected for the \"central-engine\" of the two types of progenitors (accretion timescale below 0.1 s for a compact binary merger, envelope infall-time of 100 s for a massive star). It is more likely that the burst duration is determined by the duration of the ejection of a relativistic jet by the BH left by the core collapse and/or by the timescale for dissipating the jet kinetic energy (to produce the burst emission), both of which can be shorter than the envelope infall time. That should certainly be the case if the core-collapse of massive stars accounts for the long-bursts that last only a few seconds, suggesting that the death of massive stars could also produce bursts lasting as short as 0.5 s, as for GRB 090510.  Zhang et al (2009) have proposed a replacement of the above short-long GRB dichotomy with one based on the properties expected for each type of progenitor: type I bursts are GRBs without associated supernovae, occurring in hosts with a low star-formation rate or having large offsets from the host center, while type II bursts are GRBs occurring in hosts with high star-formation rate, tracking the host light, and with associated supernovae or without deep limits on it. While recognizing that type I bursts are more likely to originate from binary mergers and type II bursts from massive stars, Zhang et al (2009) also suggest that some (or most) of the higher-luminosity short bursts (like GRB 090510) are of type II and originate from massive stars. The possibility that GRB 090510 is a short burst arising from a massive stellar progenitor also finds some support in the work of Virgili et al (2011), who argue that the luminosity-redshift distribution and the peak flux distribution of short GRBs can be explained only if a large fraction of them follow the star-formation history and are of type II."
    },
    "1005/1005.3923_arXiv.txt": {
        "abstract": "We find a unique direction in the CMB sky around which giant rings have an anomalous mean temperature profile. This direction is in very close alignment with the afore measured anomalously large bulk flow direction. Using Monte Carlo simulations, we estimate the significance of the giant rings at the $3\\sigma$ level and the alignment with the bulk flow at $2.5\\sigma$. We argue that a cosmic defect seeded by a pre-inflationary particle could explain the giant rings, the large bulk flow and their alignment. ",
        "introduction": "One of the key assumptions in modern cosmology is statistical isotropy. The detailed data from the Wilkinson Microwave Anisotropy Probe (WMAP) provides an opportunity to test this assumption. Indeed many authors have studied this issue directly and indirectly using various approaches and claimed the existence of a number of anomalies in the data (see e.g. \\cite{Spergel:2003, Copi:2007, Hajian:2007pi, Bennett:1992, Hinshaw:1996, Tegmark:2003, Schwarz:2004, Land:2005, Vielva:2004, Mukherjee:2004, Cruz:2006apj, Cruz:2006mnras, Cruz:2007, Cruz:2008, Eriksen:2004, Hansen:2004, Eriksen:2007, Hoftuft:2009, Groeneboom:2009, Joshi:2009mj, Souradeep:2006dz} and  \\cite{Abramo:2010gk, Copi:2010na} for   recent reviews).  In this paper we propose another approach to test statistical isotropy. We study how much giant rings in the Cosmic Microwave Background (CMB) sky deviate from random behavior and estimate the significance of the deviation. In section II we define the {\\it rings score} as a function of the direction the giant rings surround and generate a rings score map for the masked Internal Linear Combination (ILC) map. We show that the ILC rings score map has a clear peak. We estimate the giant rings in the ILC map as a 3$\\sigma$ deviation from $\\Lambda$CDM. Moreover, we find that the giant rings are aligned with another reported $\\Lambda$CDM anomaly \\cite{Watkins, BulkFlow, Lavaux} in the form of a large cosmic bulk flow.  In section III we discuss a cosmological scenario that could explain the giant rings, the large bulk flow and their alignment.  It is this cosmological scenario, which  involves the effects of a pre-inflationary particle \\cite{Sunny, PIP}, that actually motivated us to look for these giant rings in the first place. Section IV is devoted to discussion. ",
        "conclusions": "In this paper we reported on a novel unexpected feature of the CMB sky -- giant anomalous rings. The significance of these giant rings by themselves, much like other ``anomalous\" features of the CMB, is far from being overwhelming as it is merely a $3\\sigma$ effect. Moreover, much like the absence of large angular correlations \\cite{Spergel:2003, Copi:2007}, our findings are weaker when considering the full ILC map with no mask -- the peak in the rings score map remains aligned with the bulk flow direction, but the score is no longer significant vs. random maps. This could either mean that this is due to the contamination of the unmasked data or that the feature is weaker than suggested by the masked map. With data from Planck we should be able to determine which possibility is right. In fact, since Planck is expected to clean much of the Galactic foreground, we should be able to include the small $\\theta$ data as well and see if the signal increases.  Estimating the significance of features in the data is tricky in general and in the case of statistical isotropy in particular. Indeed many of the reported large scale anomalies in the CMB that imply violation of statistical isotropy were recently deemed as stemming from a-posteriori choices of estimators \\cite{WMAP7, Efstathiou:2009, Groeneboom:2009apj, Hanson:2009, Zhang:2009} that amplified the significance of the results. Among the claims against these anomalies is that they surfaced from a search of oddities in the data with no independent experimental evidence or prior theoretical motivation.  The giant rings are different  in this regard. First, the search for them was motivated by a theoretical scenario which by construction violates statistical isotropy. Secondly, they are aligned with another large scale {\\it non } CMB ``anomaly\"  -- the bulk flow. This increases, in our opinion, the chance that our findings could eventually lead to a real challenge for statistical isotropy. For this to happen more data is needed.  Luckily there are two clear predictions of our scenario that could be tested already with Planck data. First, the weak gravitational lensing of the CMB by the SSCD (assuming it has the magnitude required to produce the large bulk flow) is quite distinct \\cite{ab} (again because of the long range gravitational potential it induces) and should be detected by Planck. Secondly, the measurement of peculiar velocities via the  kinetic Sunyaev-Zel'dovich effect should improve quite significantly with Planck. This should enable testing the claims of \\cite{Kashlinsky}, which are based on data from WMAP, and determine if the bulk flow is indeed anomalous at even larger scales (which will increase the significance dramatically), and in which direction it points. If our scenario is correct then as one increases the size of the survey, the usual $\\Lambda$CDM effects should become more negligible compared to the SSCD effect and the measured bulk flow will be more aligned with the giant rings."
    },
    "1005/1005.1926_arXiv.txt": {
        "abstract": "We propose a new technique for weak gravitational lensing in the radio band making use of polarization information. Since the orientation of a galaxy's polarized emission is both unaffected by lensing and is related to the galaxy's intrinsic orientation, it effectively provides information on the unlensed galaxy position angle. We derive a new weak lensing estimator which exploits this effect and makes full use of both the observed galaxy shapes and the estimates of the intrinsic position angles as provided by polarization. Our method has the potential to both reduce the effects of shot noise, and to reduce to negligible levels, in a model-independent way, all effects of intrinsic galaxy alignments. We test our technique on simulated weak lensing skies, including an intrinsic alignment contaminant consistent with recent observations, in three overlapping redshift bins. Adopting a standard weak lensing analysis and ignoring intrinsic alignments results in biases of 5-10\\% in the recovered power spectra and cosmological parameters. Applying our new estimator to one tenth the number of galaxies used for the standard case, we recover both power spectra and the input cosmology with similar precision as compared to the standard case and with negligible residual bias, even in the presence of a substantial (astrophysical) scatter in the relationship between the observed orientation of the polarized emission and the intrinsic orientation. Assuming a reasonable polarization fraction for star-forming galaxies, and no cosmological conspiracy in the relationship between polarization direction and intrinsic morphology, our estimator should prove a valuable tool for weak lensing analyses of forthcoming radio surveys, in particular, deep wide field surveys with e-MERLIN, MeerKAT and ASKAP and ultimately, definitive radio lensing surveys with the SKA. ",
        "introduction": "\\label{sec:intro} The bending of light by mass inhomogeneities in the Universe results in coherent distortions in the observed images of faint background galaxies. In recent times, significant progress has been achieved in measuring this weak gravitational lensing effect (see e.g.\\ \\citealt{massey10} for a recent review). Weak lensing on cosmological scales (or `cosmic shear') was first detected only ten years ago \\citep{wittman00, kaiser00, bacon00, vanwaerbeke00}, and has since been measured with steadily improving precision by a large number of groups (see e.g.\\ \\citealt{fu08} for the most recent constraints from the CFHT Legacy Survey). Since the underlying physics of weak lensing is clean, future cosmic shear measurements that include distance information on the source galaxies, are considered to be one of the most promising techniques for constraining the growth of cosmological fluctuations over cosmic time. Cosmic shear therefore offers great promise for constraining a variety of cosmological parameters including the amplitude of the fluctuations, neutrino masses and the properties of dark energy (e.g.\\ \\citealt{albrecht06, peacock06}).  In order to measure weak lensing, one must average over the observed shapes of a sufficient number of background galaxies --- since one does not know, a priori, the intrinsic shape of any one galaxy, the lensing distortion cannot be recovered from a single galaxy image. If one assumes that there is no correlation in the intrinsic shapes of galaxies, then by averaging over a sufficient number of source galaxies, one obtains an unbiased estimate of the distortion induced by lensing.  There are two caveats to this technique. First, there is both theoretical motivation (e.g.\\ \\citealt{crittenden01, catelan01, mackey02, jing02, hirata04}) and observational evidence \\citep{brown02, heymans04, mandelbaum06, mandelbaum09, hirata07, brainerd09} for believing that the intrinsic shapes of galaxies are not completely random. Such intrinsic shape correlations will mimic a lensing signal, introducing spurious power (the so-called intrinsic-intrinsic or `II' term) into estimates of the auto-correlations of the lensing distortion (or `shear') field. Moreover, intrinsic correlations can introduce spurious anti-correlations between the shear estimated from galaxies which are widely separated in redshift (the so-called lensing-intrinsic interference or `GI' term; \\citealt{hirata04}). Of these two effects, the GI term is potentially the most worrisome as it is not easily removed from cosmic shear estimators.  Secondly, even if the intrinsic shapes of galaxies are truly random, the dispersion in galaxy shapes introduces a shot noise term. Since this dispersion is very much larger than the sought-after weak lensing signal, one requires large numbers of background galaxies in order to beat down this noise. For this reason, the vast majority of weak lensing surveys to date have been conducted in the optical. In particular, weak lensing in the radio band has lagged behind somewhat due to the much smaller galaxy number density achieved in radio surveys. However, the next generation of radio surveys such as the imminent e-MERLIN\\footnote{http://www.merlin.ac.uk/e-merlin} and LOFAR \\citep{morganti10} experiments, the SKA pathfinders, MeerKAT \\citep{booth09} and ASKAP \\citep{johnston08}, and ultimately, the SKA\\footnote{http://www.skatelescope.org} itself, will be of sufficient sensitivity to achieve a comparable source galaxy number density to planned optical surveys.  The only significant measurement of cosmic shear in the radio band is the work of \\cite{chang04} who made a statistical detection in the VLA FIRST survey. Recently, a further attempt to measure a lensing signal was applied to data from the VLA and MERLIN \\citep{patel10}. This latter work did not detect a significant lensing signal precisely because of the small number density of galaxies typically found in radio surveys. However, it was able to assess the feasibility of doing so and also proposed that systematic effects could be removed by observing the same patch of the sky in the radio and optical wavebands.  One aspect of radio surveys that is potentially useful for weak lensing is the additional polarization information which often comes for free in these surveys. In particular, all of the forthcoming surveys mentioned above will include full polarization information. Numerous authors have demonstrated that, in almost all cases of astrophysical interest, the orientation of the polarized emission from a source galaxy is unaffected by gravitational lensing \\citep{kronberg91, dyer92, faraoni93, surpi99, sereno05}. It follows that if there exists a correlation (or anti-correlation) between the intrinsic morphological orientation of a source and its polarized emission, then the observed polarization provides information on the unlensed source orientation. Such a relationship certainly exists for quasars where the polarization is closely aligned with the radio jets and this effect has already been exploited to measure gravitational lensing using polarization observations of quasars \\citep{kronberg91, kronberg96, burns04}. The number counts of future deep radio surveys will be dominated by quiescent star-forming galaxies rather than AGN. Observations in the local Universe indicate that the orientation of the polarized emission from these sources is also strongly anti-correlated with the orientation of the galaxy (e.g.\\ \\citealt{stil09}).  In this paper, we construct a new weak lensing estimator which folds in the extra information on the intrinsic orientation of galaxies which can potentially be provided by radio polarization observations. We demonstrate that, depending on the strength of the correlation between the polarized emission and the galaxy orientation, our proposed technique has the potential to reduce the impact of shot noise and to mitigate intrinsic alignment (IA) effects. Focusing on the potential impact for cosmic shear measurements, we test our technique on simulated weak lensing skies which include an IA contaminant. We demonstrate that IA biases can be reduced to negligible levels, in a model-independent way, with negligible loss in cosmological information as compared to a standard weak lensing analysis which uses a factor ten more galaxies.  Note that our proposed technique is markedly different to other techniques for mitigating the effects of IA which are generally based on nulling the contaminating signals \\citep{king02, heymans03, takada04, joachimi08, joachimi09a} or on fitting parametrized models of the IA signals (e.g.\\ \\citealt{king03, king05, bridle07}). Some of the nulling techniques are focused on removing the II term and some have been constructed to explicitly remove the GI term, while the modeling approaches attempt to model out all contaminating signals. In general, the nulling techniques are lossy in the sense that they discard useful cosmological information along with the IA contamination. The modeling techniques are less lossy but are dependent on our highly uncertain knowledge of the physics underlying IA effects. This model dependence can be mitigated to some degree by including empirical constraints on the IA signal, for example, by considering measurements of the cross correlation between the galaxy number density and observed shear fields \\citep{zhang08, bernstein09, joachimi09b, kirk10}.  All of these existing techniques remove or model the IA contamination at the two-point level (i.e.\\ in the correlation functions or power spectra). Since the technique proposed here works directly on the shear field, it separates the IA and lensing signals at the map level. It therefore does not discriminate between II and GI contamination and can, in principle, be used to make maps (albeit noisy ones) of the IA signal as well as to directly correct the lensing shear field(s) for the effects of IA. In addition to cosmic shear, in principle, our technique can also be used to correct for IA contamination in other weak lensing applications, e.g.\\ in lensing reconstructions of dark matter distributions around clusters and superclusters.  The paper is organized as follows. In Section~\\ref{sec:pol_morph}, we review the observational evidence for a correlation between the orientation of the polarized emission and the intrinsic position angle of galaxies. In Section~\\ref{sec:estimator}, we derive our new estimator and examine some of its properties. In Section~\\ref{sec:sims} we describe our simulations of weak lensing skies including an IA contaminant. The analysis of the simulations is presented in Section~\\ref{sec:analysis}. We finish by discussing our results and their implications for future radio lensing surveys in Section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  Future deep, wide field radio surveys will present a significant opportunity to measure weak lensing in the radio band with high precision for the first time. For such precise measurements, the effects of intrinsic alignments can no longer be ignored. In this paper, we have presented a technique to mitigate against IA contaminants in radio weak lensing analyses which makes use of the fact that the orientation of the polarized emission from background sources is both unaffected by gravitational lensing, and is related to the intrinsic morphological orientation of the source.  We have demonstrated the technique on simulated weak lensing skies including a contaminating IA signal. The results of these tests suggest that our method has the potential to both reduce shot noise and to strongly reject IA contamination, in a model-independent way, with negligible loss in cosmological information. We note that there are a number of limitations to the simulations presented in this paper. In particular, we have approximated the lensing shear signal as being constant for all galaxies within each pixel of each of our three redshift bins. Of course, in reality, the shear undergone by each galaxy will depend on both its angular position within the pixel and on its redshift. This is, however, not an obstacle for our proposed technique. In the absence of IA signals, it is easy to show that our new estimator for the shear in each pixel simply returns a weighted average of the shears sampled at each galaxy position when those shears are different. This average shear is, of course, the cosmological observable of interest --- the effects of sub-pixel gradients in the true shear field and the dependence on redshift of the average shear field are easily accounted for when generating model predictions for comparison with observations. In the presence of both a scatter in the intrinsic position angle estimates and a non-zero IA signal, there remains a residual bias in the estimator. However, we have already demonstrated that this bias is small and, if necessary, could be corrected for.  Although we have presented our technique in the context of weak lensing in the radio band, an extension to other wavebands, and in particular to the optical, is in principle possible and would be highly desirable. In this regard, \\cite{audit99} have already suggested that the orientation of the optical polarized emission from a galaxy also provides a proxy for the galaxy's intrinsic orientation. However, the underlying mechanism for the generation of the polarization is quite different between the radio and optical bands --- in the former case, the dominant source of polarization is synchrotron radiation whereas in the latter case, scattering by dust can also be an appreciable source. In addition, sensitive polarization measurements are more difficult, from an instrumental point of view, in the optical band.  Our proposed technique is clearly dependent on the true nature of the polarized emission from the faint star forming galaxies which will dominate the number counts of future radio surveys. In particular, the typical polarization fraction of such galaxies at $\\mu$Jy flux densities and the degree to which the orientation of the polarized emission is aligned with the morphological orientation are currently very uncertain. In addition, to apply our analysis to real data, we will need to correct the observed polarization for the rotation of the plane of polarization due to intervening cosmological magnetic fields and the effects of Faraday rotation internal to the sources themselves. These corrections should be possible by extracting the rotation measures for the sources using multi-frequency information. We hope to investigate many of these practical issues and perform the first application of our proposed technique on powerful new radio datasets from the imminent e-MERLIN and MeerKAT instruments in the not-too-distant future."
    },
    "1005/1005.0920_arXiv.txt": {
        "abstract": "We build a simple, {\\it top-down} model for the gas density and temperature profiles for galaxy clusters. The gas is assumed to be in hydrostatic equilibrium along with a component of non-thermal pressure taken from simulations and the gas fraction approaches the cosmic mean value only at the virial radius or beyond. The free parameters of the model are the slope and normalisation of the concentration-mass relation, the gas polytropic index, and slope and normalisation of the mass-temperature relation. These parameters can be fixed from X-Ray and lensing observations. We compare our gas pressure profiles to the recently proposed `Universal' pressure profile by \\cite{Arnaud09} and find very good agreement. We find that the Sunyaev-Zel'dovich Effect (SZE) scaling relations between the integrated SZE flux, $Y$, the cluster gas temperature, $T_{\\rm sl}$, the cluster mass, $M_{\\rm tot}$, and the gas mass, $M_{\\rm gas}$  are in excellent agreement with the recently observed $r_{2500}$ SZE scaling relations by \\cite{Bonamente08} and $r_{500}$ relation by \\cite{Arnaud09}. The gas mass fraction increases with cluster mass and is  given by $f_{\\rm gas}(r_{500}) = 0.1324 + 0.0284 \\,\\rm{log}\\, (\\frac{M_{500}}{10^{15}h^{-1}M_\\odot})$. This is within $10\\%$ of observed $f_{\\rm gas}(r_{500})$. The consistency between the global properties of clusters detected in X-Rays and in SZE shows that we are looking at a common population of clusters as a whole,  and there is no deficit of SZE flux relative to expectations from \\xr scaling properties. Thus, it makes it easier to compare and cross-calibrate clusters from upcoming \\xr and SZE surveys. ",
        "introduction": "\\label{sec:intro}  Large yield SZE cluster surveys promise to do precision cosmology once cluster mass-observable scaling relations are reliably calibrated. This can be done through cluster observations \\citep{Benson04, Bonamente08},  simulations \\citep{Silva04,Bonaldi07} and analytic modeling \\citep{Bulbul10}. It is well known that different astrophysical processes influence the cluster mass-observable relations non-trivially  (for example, see \\citet{Balogh01,Borgani04, Kravtsov05, Puchwein08}) which can lead to biases in determining cosmology with clusters. Alternatively, one can  `self-calibrate' the uncertainties \\citep{ 2003ApJ...585..603M,2004ApJ...613...41M, 2004PhRvD..70d3504L}.   Simplistic modeling of the intra-cluster medium (ICM), like the `isothermal $\\beta$-model' can  give rise to inaccuracies. More complex modeling needs additional assumptions (such as gas following dark matter at large radii \\citep{Komatsu01}, hereafter KS) or inclusion of less understood baryonic physics \\citep{Ostriker05}.    To partially circumvent our incomplete knowledge of cluster gas physics, we build a {\\it top-down} phenomenological model of cluster structure, taking clues from both observations and simulations. It stands on three simple, well motivated, assumptions: (i) present \\xr observations can give reliable cluster mass-temperature relations at $r<r_{500}$ which is used to calibrate our models; (ii) the gas mass fraction, $f_{\\rm gas}$, increases with radius as seen in observations \\citep {Vikhlinin06, Sun08} and in simulations \\citep{Ettori06}, with non-gravitational processes pushing the gas outwards.  It reaches values close to universal baryon fraction  at or beyond the virial radius; and (iii) there is a component of non-thermal pressure support whose value relative to thermal pressure can be inferred from biases in mass estimates found in simulations (see \\cite{Rasia04}). This simple model can reproduce the `Universal' pressure profile \\citep{Arnaud09}, \\xr gas fraction,  and SZE scaling relations in excellent agreement with observations. ",
        "conclusions": "We have constructed a {\\it top-down} model for galaxy clusters, normalized to the mass-temperature relation from \\xr observations.  The gas density and temperature profiles are found by iteratively solving the gas dynamical equation having both thermal and non-thermal pressure support. The form of the  non-thermal pressure used is taken from \\cite{Rasia04}. In our model, $f_{\\rm gas}$ becomes 0.9 ($\\frac{\\Omega_b}{\\Omega_M}$) at the cluster boundary , whereas gas is pushed  out of the cluster cores to give $f_{\\rm gas}(r_{500}) = 0.1324 + 0.0284 $ $\\rm{log} (\\frac{M_{500}}{10^{15}h^{-1}M_\\odot})$,  similar to \\xr observations.  At $r_{2500}$, the SZE scaling relations between SZE flux $Y$ and the cluster average temperature, $T_{\\rm sl}$, gas mass, $M_{\\rm gas}$, and total mass, $M_{\\rm tot}$,  show very good agreement and are within $1-\\sigma$ to the best fit line to the \\cite{Bonamente08} data. Especially, for the $Y-M_{\\rm gas}$ relation the agreement is excellent. In comparison, we also show that the Komatsu-Seljak model is in less agreement to the SZE scaling relations, especially for $Y-T_{\\rm sl}$ and $Y-M_{\\rm gas}$. Our $r_{2500}$ scaling relations can be compared to those obtained from simulations. For example, the \\cite{Nagai06}(see their table 3 for scaling parameters) radiative simulation prediction for the $Y-M_{\\rm gas}$ relation gives a $\\Delta\\chi^2 = 4.7$  w.r.t.  to the best fit \\cite{Bonamente08} relation. Recently \\cite{Bode09} have predicted SZ scalings from a mixture of N-body simulations plus semi-analytic gas models, normalized to \\xr observations for low-$z$ clusters.  The $\\Delta\\chi^2$ of their model is 0.06 and agrees well with our results.  Further out, at $r_{500}$, the $Y-M$ scaling relation obtained for our models agree very well with those obtained from X-Ray observations \\citep{Arnaud09}. Most assuringly, the gas pressure profile in our simple phenomenological model of clusters, comes out to be within $\\sim20\\%$ beyond .1 $R_{500}$ to the observed `Universal' pressure profile given by \\cite{Arnaud09}.   Most importantly, the fact that \\xr normalised models can reproduce SZE scaling relations  well is reassuring for cluster studies. It shows that we are looking at a common population of clusters as a whole,  and there is no deficit of SZE flux relative to expectations from \\xr scaling properties. Thus, one can compare and cross-calibrate clusters from upcoming \\xr and SZE surveys with increased confidence. It also gives us confidence to extrapolate our models to larger radii in order to construct the $Y-M_{200}$ scaling relation and SZ power spectrum templates."
    },
    "1005/1005.2787_arXiv.txt": {
        "abstract": "{ We investigate the dust associated with the supernova remnant (SNR) N49 in the Large Magellanic Cloud (LMC) as observed with the {\\it Herschel} Space Observatory. N49 is unusually bright because of an interaction with a molecular cloud along its eastern edge. We have used PACS and SPIRE to measure the far IR flux densities of the entire SNR and of a bright region on the eastern edge of the SNR where the SNR shock is encountering the molecular cloud. Using these fluxes supplemented with archival data at shorter wavelengths, we estimate the dust mass associated with N49 to be about 10 $M_{\\odot}$. The bulk of the dust in our simple two-component model has a temperature of 20-30 K, similar to that of nearby molecular clouds. Unfortunately, as a result of the limited angular resolution of {\\it Herschel} at the wavelengths sampled with SPIRE, the uncertainties are fairly large.  Assuming this estimate of the dust mass associated with the SNR is approximately correct, it is probable that most of the dust in the SNR arises from regions where the shock speed is too low to produce significant X-ray emission. The total amount of warm 50-60 K dust is $\\sim$0.1 or 0.4 $M_{\\odot}$, depending on whether the dust is modeled in terms of carbonaceous or silicate grains. This provides a firm lower limit to the amount of shock heated dust in N49. ",
        "introduction": "N49 is a bright X-ray and optical supernova remnant (SNR) in the Large Magellanic Cloud (Long et al. 1981), associated with a soft gamma ray repeater (Cline et al 1982) and unresolved X-ray point source (Rothschild et al. 1994). It has an asymmetric surface-brightness distribution at essentially all wavelengths, which is due to an interaction with a molecular cloud that is located on the southeast limb of the SNR (Vancura et al. 1992; Banas et al 1997). Based on a simple Sedov model for the SNR, Hughes et al. (1998) concluded that the swept-up mass was the order of 200 $M_{\\odot}$. Park et al. (2003) recently conducted a detailed study of its X-ray properties using {\\it Chandra}, confirming that the SNR is relatively young, on the order of 6\\,000 yrs.  They find that although most of the emission is dominated by interstellar material (ISM), the {\\it Chandra} observations appear to show ``explosion fragments'' on the SW side of the SNR in X-rays. According to Bilikova et al. (2007), the X-ray temperature ($\\sim$1.2$\\times$10$^{7}$ K) implies a shock velocity of 920 km s$^{-1}$, while the optical filaments which reflect shocks traversing denser gas have typical line-of-sight velocities of 250 km s$^{-1}$. Hill et al. (1995) estimated the mass of the progenitor of N49 to be about 20 $M_{\\odot}$ based on the association with the 10 Myr old association LH 53. A 1720 MHz OH maser emission comes from a region toward the southwest. N49 is almost certainly the result of a Type II explosion.     Apparently first detected at mid-IR wavelengths with {\\it IRAS} (Graham et al.\\ 1987), the first mid-IR images of the SNR were obtained with {\\it Spitzer} (Williams, R. et al. 2006), where N49 is visible not only in the MIPS 24 and 70 $\\mu$m bands, but also in the IRAC 3.6, 4.5, 5.6 and 8 $\\mu$m bands. IRS spectra of the eastern limb show strong line emission from O, Ne, Ar, Si, S, and Fe and H$_{2}$ in the 5-37 $\\mu$m band (Williams, R. et al 2006). While lines and dust emission compete at the short wavelengths, van Loon et al (2010) showed with the low resolution MIPS spectra of the bright blob on the eastern limb that dust is the dominant source of emission at 70 $\\mu$m, and estimated a dust mass for N49 of $\\sim$0.2 $M_{\\odot}$.    Dust in SNRs can arise either from the SN itself or from the ISM. Ejecta dust is thought to play a key role as a coolant in the formation of high-redshift galaxies (see, e.g. Morgan \\& Edwards 2003). However, given the amount of swept-up material in N49, most of the dust in N49 is expected to be of interstellar origin.    Here we describe observations of N49 obtained with the {\\it Herschel} Space Observatory. Our primary purpose is to explore the properties of dust in N49, and to obtain a better estimate of the dust mass than was heretofore possible by probing the longer wavelength bands, where larger dust grains are expected to be more apparent. Combining {\\it Herschel} data and data obtained from other observatories, we construct spectral energy distributions (SEDs) of the entire SNR and of a bright blob on the eastern limb, where the SNR is interacting with a molecular cloud. (see, e.g. van Loon et al. 2010).  We use a simple two-temperature model of the SED to estimate dust masses and temperatures. ",
        "conclusions": "We estimated the dust mass with the line-contribution-subtracted data and its temperature through SED fitting. First, we estimated the contribution of the synchrotron emission to the mid- to far-IR region. Dickel \\& Milne (1998) measured the flux densities at 3.5/6/13/20 cm for the whole SNR. By extrapolating the flux densities at this radio wavelength to the mid-IR region, we found that the contribution is negligibly small ($<$2 $\\%$) to the 24/70/100/160/250 $\\mu$m flux densities. In the SPIRE 350/500 $\\mu$m bands, the contribution is large; $\\sim$14 $\\%$ for 350 $\\mu$m and $\\sim$79 $\\%$ for 500 $\\mu$m. Before fitting, we subtracted the contribution of the synchrotron emission to the 350 $\\mu$m flux density for the whole SNR and E blob, assuming that the contribution can be also applied to the E blob. In the fittings we excluded 500 $\\mu$m data.   When we assume optically thin thermal radiation at a dust temperature $T_{d,i}$ of a component $i$ and a single dust species with the radius $a$, the observed flux density $F_{\\nu}$ is written by       \\begin{equation} F_{\\nu} = \\left(\\frac{4}{3}a\\rho D^{2}\\right)^{-1} \\times \\sum_{i} M_{d,i}B_{\\nu}(T_{d,i})Q_{\\nu,i}, \\label{dust} \\end{equation}    \\noindent where $M_{d,i}$ is the dust mass of the component $i$, $B_{\\nu}$($T_{d,i}$) is the Planck function of $T_{d,i}$, and $Q_{\\nu,i}$ is the absorption efficiency, $\\rho$ is the dust grain density and $D$ is a distance to the observed dust from us (Kwok 2007). Here, we adopted $D$=50 kpc (Schaefer 2008).  We considered two types of grains, amorphous carbon and silicates.  For amorphous carbon the optical constants were taken from Zubko et al. (1996). We adopted $\\rho$=2.26 g cm$^{-3}$. To simplify the model, we assume $a$ = 0.1$\\mu$m uniform grain radius. Next, we considered astronomical silicate only. The optical constants were taken from Draine \\& Lee (1984) and a $\\rho$ of 3.6 g cm$^{-3}$ was adopted. Due to the lower absorption efficiency at $>$70 $\\mu$m than amorphous carbon, a larger dust mass would be estimated. The dust mass estimated for the silicate-only case would be an upper limit. In the fittings we found the best fit $M_{d,i}$ and $T_{d,i}$ to the observations by least-square technique. In this process, $M_{d,i}$ and $T_{d,i}$ are the free parameters, and all the data are equal of weight.    We first considered a single temperature model to fit the SED. For the whole SNR, the best fit had a $T_{d}$ of 33 K and 34 K for amorphous carbon and for silicate respectively, and a derived $M_{d}$ of 3.6 $M_{\\odot}$ and 4.6 $M_{\\odot}$. For the E blob, $T_{d}$ is 35 K (carbon case) and 39 K (silicate case), and $M_{d}$ is 0.47 $M_{\\odot}$ and 0.60 $M_{\\odot}$. The temperature is fairly similar to the temperature derived by Graham et al.\\ (1987) from {\\it IRAS} data. However, we found that a better fit could be obtained if we considered a model with two-temperatures instead of one. The one-component fittings cannot simultaneously explain the $<$24 $\\mu$m and $>$160 $\\mu$m flux. The implied dust masses and temperatures for the two-temperate model are summarized in Table 2 for the two types of grains.     The best-fit two-temperature model, assuming silicate grains, for the SEDs is shown in Fig 2.  The warm component has a temperature of $\\sim$60 K, and is surely due to dust in shocked gas, while the cold component has a temperature of $\\sim$20 K, similar to that found in molecular clouds (16-24 K, Bernard et al. 2008; Gordon et al. in 2010). The contribution from the cold component to the total flux in the {\\it Herschel} bands is $>$80 $\\%$. The dust mass and temperature of the  warm component in the E blob are very close to van Loon et al. (2010; 0.2 $M_{\\odot}$ and 43 K). The total mass of dust associated with N49 is about 10 $M_{\\odot}$ and is dominated by dust at cold temperatures, dust which was difficult to observe with {\\it Spitzer}. It is tempting therefore to identify the warm component with dust in the X-ray plasma and the cool component not with the SNR, but with the molecular cloud.  However, this assignment is very uncertain, because dust temperature variations are also expected as a result of the large variations of post-shock gas density that are known to exist in N49.   To place SNR N49's dust mass estimate into the context of SNR development, we plot the dust mass vs. age for all LMC and SMC SNRs studied to data, as noted in Fig 3. Note that except for N49 the dust masses are estimated based on $<$70 $\\mu$m data. The relation can be with the free expansion phase early on, followed by a Sedov phase (radius $\\propto$ (age)$^{2/5}$; cf. Badenes et al. 2010).    The mass of the warm dust component in N49 is 0.1 $M_{\\odot}$ assuming amorphous carbon and 0.4 $M_{\\odot}$  assuming silicates.  Adopting a dust-to-gas mass ratio of $\\sim$250 (Meixner et al. 2010) and assuming that only the warm dust is associated with shocked gas, the swept-up mass of the SNR ranges from 25 to 100 $M_{\\odot}$. This is comparable to the mass estimated from X-ray observations, 200 $M_{\\odot}$ (Hughes et al. 1998), assuming the true grain characteristics are similar to those of silicates. A much higher  (25$\\times$ larger)  swept-up mass is estimated if we included the cold dust in our estimate, which suggests that in the context of our simple model there is a significant molecular cloud component to the cold dust we associate with the SNR, and that the shocks propagating into the molecular cloud may be too slow to produce X-rays.  While more sophisticated modeling of N49 may refine the mass estimates somewhat, the basic limitation in arriving at a more precise picture of the nature of the interaction with the molecular cloud and the amount of shocked molecular gas in N49 is the limited angular resolution of {\\it Herschel} at longer wavelengths.   \\begin{figure} \\centering \\includegraphics[width=7.0cm]{f3.eps} \\caption{The relation between the dust mass and age. The data are taken from Seok et al. (2008) for 0509-67.5, 0519-69.0, N132D, and N49B; Williams, B. et al. (2008); Sandstrom et al. (2009) and Rho et al. (2009) for 1E 0102.2-7219; Williams, B. et al. (2006) for N132D, N23, 0453-68.5, and N49B; van Loon et al. (2010) and present work for N49.\\label{model}}  \\end{figure}"
    },
    "1005/1005.1671_arXiv.txt": {
        "abstract": "{We report the detection of strong absorption by interstellar hydrogen fluoride along the sight-line to the submillimeter continuum source G10.6--0.4 (W31C).  We have used Herschel's HIFI instrument, in dual beam switch mode, to observe the 1232.4763~GHz $J=1-0$ HF transition in the upper sideband of the Band 5a receiver.  The resultant spectrum shows weak HF emission from G10.6--0.4 at LSR velocities in the range --10 to --3 km s$^{-1}$, accompanied by strong absorption by foreground material at LSR velocities in the range 15 to 50 km~s$^{-1}$.  The spectrum is similar to that of the 1113.3430~GHz $1_{11}-0_{00}$ transition of para-water, although at some frequencies the HF (hydrogen fluoride) optical depth clearly exceeds that of para-H$_2$O.  The optically-thick HF absorption that we have observed places a conservative lower limit of $1.6 \\times 10^{14}\\, \\rm cm^{-2}$ on the HF column density along the sight-line to G10.6--0.4.  Our lower limit on the HF abundance,  $6 \\times 10^{-9}$ relative to hydrogen nuclei, implies that hydrogen fluoride accounts for between $\\sim 30\\%$ and 100$\\%$ of the fluorine nuclei in the gas phase along this sight-line.  This observation corroborates theoretical predictions that --  because the unique thermochemistry of fluorine permits the exothermic reaction of F atoms with molecular hydrogen -- HF will be the dominant reservoir of interstellar fluorine under a wide range of conditions.} ",
        "introduction": "Fluorine is alone among the elements of the periodic table in having a neutral diatomic hydride with a larger binding energy than that of molecular hydrogen.  As a result, fluorine atoms are unique in reacting exothermically with the dominant constituent of interstellar molecular clouds, H$_2$, producing hydrogen fluoride via the reaction\\footnote{This reaction is arguably the most extensively-studied bimolecular chemical reaction, and has been the subject of numerous experimental and theoretical investigations (Zhu et al.\\ 2002 and references therein).  Although it possesses an activation energy barrier of order $500$~K, it is believed to be moderately rapid even at low temperatures as a result of quantum mechanical tunneling.} $$\\rm F + H_2 \\rightarrow HF + H. \\eqno(1)$$ Theoretical models for the interstellar chemistry of fluorine-bearing molecules (Neufeld, Wolfire \\& Schilke 2005; Neufeld \\& Wolfire 2009) have made the prediction that hydrogen fluoride will be the dominant reservoir of gas-phase fluorine over a wide range of conditions as a direct result of reaction (1).  In particular, HF (hydrogen fluoride) is expected to be abundant even close to molecular cloud surfaces that are irradiated by the interstellar ultraviolet radiation field; indeed, HF will be the dominant fluorine-bearing species wherever hydrogen is predominantly molecular.  Thus, these models make the surprising prediction that HF can be the most abundant molecule after H$_2$ close to UV-irradiated cloud surfaces; here the expected abundance of HF exceeds that of CO despite the fact that F nuclei are less abundant than C nuclei by more than four orders of magnitude (Asplund et al.\\ 2009).  Indirect support for this prediction has come from the detection of the CF$^+$ molecular ion, by means of ground-based observations of its $J=1-0$, $J=2-1$ and $J=3-2$ transitions towards the Orion bar (Neufeld et al.\\ 2006) and of its $J=1-0$ transition towards several other UV-irradiated molecular clouds (Neufeld et al.\\ 2010, in preparation).  Because the presumptive production mechanism for CF$^+$ is the reaction $$\\rm HF + C^+ \\rightarrow CF^+ + H, \\eqno(2)$$ these detections of CF$^+$ argue for a significant overlap of the regions where HF and C$^+$ are both abundant, and thus for the presence of large HF abundances close to cloud surfaces where the C$^+$/CO ratio is large.  The detection of interstellar HF itself, however, is severely hampered by atmospheric absorption.  Because HF has the smallest moment of inertia of any molecule containing a heavy element, its pure rotational transitions lie at high frequencies inaccessible from ground-based observatories.  To date, the only detection of interstellar HF has been obtained toward a single source by means of observations (Neufeld et al.\\ 1997) performed with the {\\it Infrared Space Observatory (ISO)}.  In that study, the $J=2-1$ transition at 2.463~THz was detected in absorption toward the strong submillimeter continuum source Sgr B2.  Because the spectral coverage of {\\it ISO} did not extend to the 1.232~THz $J=1-0$ transition, the search for HF in absorption was impossible except in exceptional regions where the $J=1$ state was significantly populated by radiative excitation.  The Heterodyne Instrument for the Far-Infrared (HIFI\\footnote{HIFI has been designed and built by a consortium of institutes and university departments from across Europe, Canada and the United States under the leadership of SRON Netherlands Institute for Space Research, Groningen, The Netherlands and with major contributions from Germany, France and the US. Consortium members are: Canada: CSA, U.~Waterloo; France: CESR, LAB, LERMA, IRAM; Germany: KOSMA, MPIfR, MPS; Ireland, NUI Maynooth; Italy: ASI, IFSI-INAF, Osservatorio Astrofisico di Arcetri- INAF; Netherlands: SRON, TUD; Poland: CAMK, CBK; Spain: Observatorio Astron\\'omico Nacional (IGN), Centro de Astrobiolog\\'a (CSIC-INTA). Sweden: Chalmers University of Technology - MC2, RSS \\& GARD; Onsala Space Observatory; Swedish National Space Board, Stockholm University - Stockholm Observatory; Switzerland: ETH Zurich, FHNW; USA: Caltech, JPL, NHSC.}; De Graauw et al.\\ 2010), on board the {\\it Herschel Space Observatory} (Pilbratt et al.\\ 2010), provides access to the HF $J=1-0$ transition for the first time.  Thus, absorption line spectroscopy of bright submillimeter continuum sources, performed with HIFI, can provide a sensitive probe of HF in its {\\it ground rotational state} ($J=0$).  As part of the PRISMAS (''PRobing InterStellar Molecules with Absorption line Studies'') Key Program, we will exploit this new capability by performing HIFI observations of eight strong submillimeter continuum sources with sight-lines that are known to intercept foreground molecular material.  In this {\\it Letter}, we present the first results of our search for hydrogen fluoride, obtained from observations of the sight-line to G10.6--0.4 (W31C).  G10.6--0.4 is a region of high-mass star-formation.  It is one of three bright HII regions within the W31 complex, and an extremely luminous submillimeter and infrared continuum source; Wright, Fazio \\& Low (1977) estimate its infrared luminosity as $\\sim 10^7 L_\\odot$.  Using observations of foreground HI 21 cm absorption, Fish et al.\\ (2003) have obtained a kinematic distance estimate of $4.8^{+0.4}_{-0.8}$~kpc for G10.6--0.4, a value that places the source within the so-called ``--30~km~s$^{-1}$'' spiral arm.  The sight-line to G10.6--0.4 intersects several foreground molecular clouds, the arrangement of which has been elucidated by Corbel \\& Eikenberry (2004; see their Figure 8).    Because G10.6--0.4 shows a large continuum flux at infrared and radio wavelengths, this sight-line has proven one of the most valuable in the Galaxy for the study of interstellar gas by absorption line spectroscopy.  Indeed, line absorption by foreground gas has been detected over three and one-half decades in wavelength, from species as diverse as atomic oxygen (at 63~$\\mu$m; Keene et al.\\ 1999), $^{13}$CH$^+$ ($J=1-0$ at 361~$\\mu$m; Falgarone, Phillips \\& Pearson 2005), HCO$^+$ ($J=1-0$ at 3.36~mm; Keene et al.\\ 1999; Godard et al.\\ 2010), and atomic hydrogen (at 21~cm; Fish et al.\\ 2003). ",
        "conclusions": "Our detection of optically-thick absorption in the HF $J=1-0$ transition over a broad range of LSR velocities corroborates the theoretical prediction that HF is the dominant reservoir of gas-phase fluorine under a wide variety of interstellar conditions.  This conclusion suggests that HIFI observations of HF will provide a useful probe of diffuse molecular gas throughout the Galaxy, and -- if its high abundance is confirmed along other similar sight-lines for which upcoming HIFI observations are planned -- a potentially valuable surrogate for molecular hydrogen. Along sight-lines where absorbing material covers a substantial range of LSR velocities, HF absorption should be detectable (although spectrally unresolved) even at the resolution of Herschel's SPIRE instrument (Griffin et al.\\ 2010); this opens up the possibility of detecting hydrogen fluoride toward suitable extragalactic sources as well."
    },
    "1005/1005.4781_arXiv.txt": {
        "abstract": "*{Pulsar Wind Nebulae  (PWNe) are bubbles or relativistic plasma that form when the pulsar wind is confined by the SNR or the ISM. Recent observations have shown a richness of emission features that has driven a renewed interest in the theoretical modeling of these objects. In recent years a MHD paradigm has been developed, capable of reproducing almost all of the observed properties of PWNe, shedding new light on many old issues. Given that PWNe are perhaps the nearest systems where processes related to relativistic dynamics can be investigated with high accuracy, a reliable model of their behavior is paramount for a correct understanding of high energy astrophysics in general. I will review the present status of MHD models: what are the key ingredients, their successes, and open questions that still need further investigation.}  \\abstract{Pulsar Wind Nebulae  (PWNe) are bubbles or relativistic plasma that form when the pulsar wind is confined by the SNR or the ISM. Recent observations have shown a richness of emission features that has driven a renewed interest in the theoretical modeling of these objects. In recent years a MHD paradigm has been developed, capable of reproducing almost all of the observed properties of PWNe, shedding new light on many old issues. Given that PWNe are perhaps the nearest systems where processes related to relativistic dynamics can be investigated with high accuracy, a reliable model of their behavior is paramount for a correct understanding of high energy astrophysics in general. I will review the present status of MHD models: what are the key ingredients, their successes, and open questions that still need further investigation.} ",
        "introduction": "\\label{sec:1} When the ultra-relativistic wind from a pulsar interacts with the ambient medium, either the SNR or the ISM, a bubble of non-thermal relativistic particles and magnetic field, known as Pulsar Wind Nebula or ``Plerion''  (PWN), is formed. The Crab Nebula is undoubtedly the best example of a PWN, and it is often considered the prototype of this entire class of objects, to the point that models of PWNe are, to a large extent, based on what is known in this single case. The first theoretical model of the structure and the dynamical properties of PWNe was presented by Rees \\& Gun \\cite{ree74}, further developed in more details by Kennel \\& Coroniti \\cite{ken84a,ken84b} (KC84 hereafter), and is based on a relativistic MHD description.  The MHD paradigm is based on three key assumptions: \\begin{itemize} \\item that the larmor radii of the particles is much smaller than the typical size of the nebula, and particles are simply advected with the magnetic field. This is true up to energies of order of the pulsar's voltage, where the larmor radius becomes comparable with the typical size of the system. \\item That radiative losses are negligible, or at least that they can be accounted for by renormalizing the pulsar spin-down luminosity. This again can be proved to be true in the case of Crab Nebula (and to some extent also in other systems with good spectral coverage), where the synchrotron spectrum shows that the particles carrying the bulk of the energy have a typical lifetime for synchrotron cooling longer than the age of the nebula. \\item That we are dealing with almost pure pair plasma, and dispersive or hybrid effects (separation of scales) due to the presence of heavier ions are absent. While there is no direct evidence for the absence of ions, standard pulsar wind theory, and the success of the MHD model of PWNe suggest that, from a purely dynamical point of view, there is no need for this extra component. \\end{itemize}  In it simplest form \\cite{ree74} the MHD model of PWNe can be summarized as follow (see Fig.~\\ref{fig:4}): the ultra-relativistic pulsar wind is confined inside the slowly expanding SNR, and slowed down to non relativistic speeds in a strong termination shock (TS). At the shock the plasma is heated, the toroidal magnetic field of the wind is compressed, and particles are accelerated to high energies. These high energy particles and magnetic field produce a post-shock flow which expands at a non relativistic speed toward the edge of the nebula.  Despite its simplicity the MHD model can explain many of the observed properties of PWNe, and until now no observation has been presented that could rule it out.  The presence of an under-luminous region, centered on the location of the pulsar, is interpreted as due to the ultra-relativistic unshocked wind. Polarization measures \\cite{wil72,vel85,sch79,hic90,mic91,dod03,kot06,hes08} show that emission is highly polarized and the nebular magnetic field is mostly toroidal, as one would expect from the compression of the pulsar wind, and it is consistent with the inferred symmetry axis of the system. The pressure anisotropy associated to the compressed nebular toroidal magnetic field \\cite{beg92,van03}, explains the elongated axisymmetric shape of many PWNe ({\\it i.e.} Crab Nebula, 3C58). The MHD flow from the TS to the edge of the nebula also leads to the prediction that PWNe should appear bigger at smaller frequencies: high energy X-rays emitting particles are  present only in the vicinity of the TS, having a shorter lifetime for synchrotron losses, compared to Radio particles which fill the entire volume, having negligible losses on the age of the nebula. This increase in size at smaller frequencies is observed in the Crab Nebula \\cite{ver93,bie97,ban98}.  However one must bear to mind that not all properties of PWNe can be explained within the MHD framework, which, ultimately, only provides a description of the flow dynamics. For example, the acceleration of particles at the TS that accounts for the continuous, non-thermal, very broad-band spectrum, extending from Radio to X-rays \\cite{ver93,ban99,wei00,wil01,mor04}, is usually assumed as given. The MHD model provides no hint to the reason why the injection spectrum looks like a broken powelaw, with no sign of a Maxwellian component at lower energies. Moreover the MHD description might prove faulty if applied to particles responsible for the emission in the 10-100 MeV band, whose larmor radii are comparable to the size of the TS, and can lead to wrong conclusion on their expected behavior.  The MHD model of PWNe has been used, by comparing observations with the predictions of numerical simulations, to  constrain some of the properties of the pulsar wind, at least at the distance of the TS. While it is not possible to derive the Lorentz factor of the wind, or its multiplicity, it is possible to constrain the ratio between Poynting flux and kinetic energy, the latitudinal dependence of the energy flux, and the presence of a dissipated equatorial current sheet. This shows that nebular properties can be used to derive informations on the conditions of the pulsar wind at large distances. ",
        "conclusions": ""
    },
    "1005/1005.4748_arXiv.txt": {
        "abstract": "KMOS is a multi-object near-infrared integral field spectrometer with 24 deployable pick-off arms. Data processing is inevitably complex. We discuss specific issues and requirements that must be addressed in the data reduction pipeline, the calibration, the raw and processed data formats, and the simulated data. We discuss the pipeline architecture. We focus on its modular style and show how these modules can be used to build a classical pipeline, as well as a more advanced pipeline that can account for both spectral and spatial flexure as well as variations in the OH background. A novel aspect of the pipeline is that the raw data can be reconstructed into a cube in a single step. We discuss the advantages of this and outline the way in which we have implemented it. We finish by describing how the QFitsView tool can now be used to visualise KMOS data. ",
        "introduction": "\\label{sec:overview}  KMOS\\cite{sha10a,sha10b} is a fully cryogenic seeing limited multi-object integral field spectrometer, which is being built by a consortium of German and British institutes as one of the second generation VLT instruments. It is equipped with 24 integral field units, each with a 2.8\\arcsec$\\times$2.8\\arcsec\\ field of view, that can be deployed on-the-fly using robotic arms anywhere within a 7\\arcmin\\ patrol field. The number and size of the IFUs, and the spectral resolution, have been chosen according to the key science driver, to study the morphology and kinematics of emission lines in high redshift galaxies. Each of the 3 identical segments in the instrument comprises 8 pick-off arms (including K-mirrors to keep the fields oriented correctly on sky), filter wheels, the integral field unit (that uses mirrors to slice the image and re-arrange it along a pseudo-slit), and a spectrograph with a 2k$\\times$2k HAWAII-2RG detector. A number of gratings are available, allowing observations across various bands from 0.8--2.5\\micron\\ at a resolution of $R\\sim3500$ in order to enable science observations to probe between the bright atmospheric OH lines. An additional grating enables observation of the H and K bands simultaneously at a lower resolution. Since the pixel scale is nominally 0.2\\arcsec, a single exposure with KMOS generates 336 slit spectra, each about 2000 pixels in the spectral direction and 14 pixels wide spatially. Each slit spectrum is separated from its neighbours with a gap of a few blank pixels. Although there are a large number of spectra, the absolute data volume is not excessive: in a typical night the instrument will produce only $\\sim$5\\,Gbyte of raw data (although it could be up to $\\sim$50\\,Gbyte in the extreme case of short exposures). The pipeline modules (recipes) are written in C using the ESO Common Pipeline Library. They can be executed using the standard ESO pipeline tools, either in the command line with {\\em EsoRex} or with the {\\em Gasgano} GUI.  KMOS is currently in its manufacturing phase, and has already had its technical first light\\cite{sha10a,sha10b}. The preliminary acceptance in Europe and first light on sky at the VLT in Chile are both expected during 2011. Further details on specific aspects of the instrument, including hardware, software, control, and testing, can be found elsewhere in these proceedings. ",
        "conclusions": ""
    },
    "1005/1005.3959_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{intro}  The Voyager 1 and 2 encounters with Jupiter of March and July 1979 (jovian northern fall equinox) offered the first opportunities to map Jupiter in the thermal infrared with a latitude resolution of better than 10\\dg\\ - then unobtainable from the ground - and a spectral resolution sufficient to measure the abundances of trace gas species. The tool for this mapping was IRIS \\citep[the InfraRed Interferometer and Spectrometer,][]{hanel77}, and trace gases of interest included ethane (\\ethane ) and acetylene (\\acet ) - secondary species derived from photolysis of the primary carbon-bearing molecule, methane (\\methane ). Although some preliminary findings were published in a conference report \\citep{maguire84}, a full radiative transfer model was never applied to retrieve a meridional trend in the abundances.  Two decades later, the Cassini spacecraft swung by Jupiter in a distant flyby maneuver (137 R$_{\\rm J}$) en route to Saturn. The official encounter period, during which time observations were made, lasted six months symmetric about the closest approach on December 30th 2000 (northern summer solstice). Carrying on board the Composite Infrared Spectrometer (CIRS), Cassini was able to re-map the planet in the thermal infrared, achieving a spatial resolution comparable to Voyager IRIS but with much higher maximum spectral resolution (0.48 \\cm , versus 3.9 \\cm\\ for IRIS, full width to half maximum (FWHM)). Meridional trends of stratospheric and tropospheric \\acet\\ and \\ethane\\ have been retrieved from CIRS data acquired during December 1-15 2000 prior to cloest approach, by \\citet{nixon07} (hereafter Paper I).  The abundances of the minor hydrocarbons are important because these species serve as tracers of atmospheric circulation. The photochemical lifetime of ethane in the stratosphere (3$\\times 10^{10}$~s at 5 mbar) is significantly greater than one Jovian year (3.7$\\times 10^{8}$~s), whereas the acetylene lifetime (5$\\times 10^{7}$~s at 5 mbar) is substantially shorter. Note that the solar cycle (3.4$\\times 10^{8}$~s) which affects the photochemistry almost equals the Jovian year. In the presence of dynamical motions on seasonal timescales, the meridional abundance trend in a relatively short-lived species (such as acetylene) is expected to show significant change, whereas a long-lived species (ethane) should present a more constant distribution. The logical means of testing this hypothesis is to intercompare the hydrocarbon abundances using the datasets of Voyager and Cassini, which encountered Jupiter at two different seasons: near northern fall equinox, and 1.75 jovian years later just after northern summer solstice.  Such a comparison is the objective of this paper. We begin by analyzing the Voyager spectra, retrieving for the first time with a full radiative transfer modeling approach the meridional abundance trends of ethane and acetylene in 1979. We also reanalyze the CIRS dataset, using a significantly revised spectroscopic atlas for ethane \\citep{auwera07} that was published since the first Cassini jovian hydrocarbons paper appeared (and is also used the IRIS data analysis here). The intensities of this new list are revised upwards by some 44\\% relative to the GEISA 2003 list used in Paper I, resulting in a corresponding decrease in abundances. We then proceed to compare the abundance profiles at the two epochs (northern fall in 1979 versus summer in 2000) discussing the implications and drawing conclusions. ",
        "conclusions": "\\label{conclusions}  In this paper we have analyzed two infrared spectral datasets of Jupiter: (i) the Voyager IRIS dataset, pertaining to the northern fall equinox of 1979, and (ii) the Cassini CIRS dataset, taken near the northern summer solstice of 2000. We have modeled the spectra to retrieve, firstly atmospheric temperatures through the $\\nu_4$ band of \\methane\\ at $\\sim1305$\\cm , and secondly the abundances of \\ethane\\ and \\acet\\ based on the derived temperatures. Our results extend from the upper troposphere to the upper stratosphere for temperature and \\acet\\ abundance, and cut-off slightly lower in the middle stratosphere for \\ethane . The results comprise the first comprehensive 2-D comparison of the abundances of \\acet\\ and \\ethane\\ in Jupiter's atmosphere at two different seasons. The results have been archived on-line for reference, as described in \\S \\ref{sect:results}.  Our new findings confirm, improve and extend the results of our previous paper \\citep{nixon07}. The distribution of \\acet\\ at the Cassini epoch is confirmed, whilst that of \\ethane\\ is corrected, largely due to revised spectral line intensities published in the interim \\citep{auwera07}. We have also extended the Cassini distributions to a second Jovian season (Voyager), showing important similarities and differences.  For ethane, the 2-D distribution is quite similar, although we see that the Cassini picture is more symmetric than that of Voyager, where \\ethane\\ is somewhat depleted in the north relative to the south. Also, the Cassini abundances are mostly higher overall, perhaps due to smaller solar distance. For acetylene the difference is much more dramatic, showing a uniform meridional profile at equinox, but strong depletions towards both poles at solstice. Therefore the solstice distribution seems to match the annual-mean picture that would be predicted by a 1-D photochemical model.  It is clear that Jupiter's stratosphere is indeed a complex environment, where chemistry is coupled to dynamics in ways that are poorly understood at present. Much additional modeling work is required before these seasonal changes can be confidently interpreted, and comprehensive radiative-dynamical-chemical models in two or three dimensions are urgently needed. The rewards of successful modeling of this data are large, and include the significant advancement of our understanding not only of the Jovian atmosphere, but also that of Saturn and the other outer planets as well.   \\clearpage {\\bf Acknowledgements}  The acquisition of CIRS data is the result of the collective efforts of a large number of people, including the following who worked on various aspects of CIRS science planning, instrument commanding, uplink, calibration and databasing: S.B. Calcutt, R.C. Carlson, M.H. Elliott, E. Guandique, M. Kaelberer, V.G. Kunde, E. Lellouch, A. Mamoutkine, P.J. Schinder, M.E. Segura, J.S. Tingley, and also many engineers and science planners at the Jet Propulsion Laboratory. We would like to thank R.A. West, A.J. Friedson and Y.L. Yung for helpful discussions during the preparation of this manuscript. During the research for and writing of this report, RKA and FMF were funded by the NASA Cassini Project, and CAN and MA were supported by the NASA Outer Planets Research Program. Portions of this work were carried out by the Jet Propulsion Laboratory, under contract with the National Aeronautics and Space Administration.  \\appendix   \\clearpage"
    },
    "1005/1005.1995_arXiv.txt": {
        "abstract": "% The INTEGRAL satellite has revealed a major population of supergiant High Mass X-ray Binaries in our Galaxy, revolutionizing our understanding of binary systems and their evolution. This population, constituted of a compact object orbiting around a massive and luminous supergiant star, exhibits unusual properties, either being extremely absorbed, or showing very short and intense flares.  An intensive set of multi-wavelength observations has led us to reveal their nature, and to show that these systems are wind-fed accretors, closely related to massive star-forming regions.  In this paper I describe the characteristics of these sources, showing that this newly revealed population is linked to the evolution of gamma-ray emitting massive stars with a compact companion. ",
        "introduction": "satellite}  The {\\it INTEGRAL} observatory is an ESA satellite launched on 17 October 2002 by a PROTON rocket on an excentric orbit. It is hosting 4 instruments: 2 $\\gamma$-ray coded-mask telescopes --the imager IBIS and the spectro-imager SPI, observing in the range 10 keV-10 MeV, with a resolution of $12\\amin$ and a field-of-view of $19\\adeg$-- a coded-mask telescope JEM-X (3-100 keV), and an optical telescope (OMC).   The $\\gamma$-ray sky seen by {\\it INTEGRAL} is very rich, since 499 sources have been detected by {\\it INTEGRAL}, reported in the $3^{rd}$ IBIS/ISGRI soft $\\gamma$-ray catalogue, spanning 3.5 years of observations in the 20-100 keV domain \\citep{bird:2007}.  214 sources were discovered by {\\it INTEGRAL}, while the remaining 285 were already known. Among these sources, there are 147 XRBs (representing 29\\% of the whole sample of sources detected by {\\it INTEGRAL}, called ``IGRs'' in the following), 163 AGNs (33\\%), 27 CVs (5\\%), and 20 sources of other type (4\\%): 12 SNRs, 2 globular clusters, 2 SGRs and 1 GRB. 129 objects still remain unidentified (26\\%).  XRBs are separated in 82 LMXBs and 78 HMXBs (each category represents 16\\% of IGRs). Among the HMXBs, there are 24 BeXBs and 19 sgXBs (representing respectively 31\\% and 24\\% of HMXBs).  It is interesting to follow the evolution of the ratio between BeXBs and sgXBs.  During the pre-{\\it INTEGRAL} era, HMXBs were mostly BeXBs systems.  For instance, in the catalogue of 130 HMXBs by \\cite{liu:2000}, there were 54 BeXBs and 7 sgXBs (respectively 42\\% and 5\\% of the total number of HMXBs).  Then, the situation changed with the first HMXBs identified by {\\it INTEGRAL}: in the catalogue of 114 HMXBs (+128 in Magellanic Clouds) of \\cite{liu:2006}, there were 60\\% of BeXBs and 32\\% of sgXBs firmly identified.  Therefore, while the ratio of BeXBs/HMXBs increased by a factor of 1.5 only, the one of sgXBs/HMXBs increased by a factor of 6. ",
        "conclusions": "The {\\it INTEGRAL} satellite has tripled the total number of Galactic sgXBs, constituted of a NS orbiting around a supergiant star. Most of these new sources are slow and absorbed X-ray pulsars, exhibiting a large $\\nh$ and long $P_\\mathrm{spin}$ ($\\sim$1ks).  The influence of the local absorbing matter on periodic modulations is different for sgOB or BeXBs, segregated in different parts of $\\nh$-$\\Porb$ or $\\nh$-$P_\\mathrm{spin}$. {\\it INTEGRAL} revealed 2 new types of sources.  First, the SFXTs, exhibiting short and strong X-ray flares, with a peak flux of 1 Crab during 1--100s, every $\\sim 100$\\,days. These flares can be explained by accretion through clumpy winds.  Second, the obscured HMXBs are persistent X-ray sources composed of supergiant stellar companions exhibiting a strong intrinsic absorption and long $P_\\mathrm{spin}$.  The NS is deeply embedded in the dense stellar wind, forming a dust cocoon enshrouding the whole binary system.  These results show the existence in our Galaxy of a dominant population of a previously rare class of high-energy binary systems: supergiant HMXBs, some exhibiting a high intrinsic absorption (\\citeauthor{chaty:2008} \\citeyear{chaty:2008}; \\citeauthor{rahoui:2008} \\citeyear{rahoui:2008}). Studying this population will provide a better understanding of the formation and evolution of short-living HMXBs.  Furthermore, stellar population models now have to take these objects into account, to assess a realistic number of high-energy binary systems in our Galaxy."
    },
    "1005/1005.4324_arXiv.txt": {
        "abstract": "The deconfinement phase transition which happens in the interior of neutron stars are investigated. Coupled with the spin evolution of the stars, the effect of entropy production and deconfinement heat generation during the deconfinement phase transition in the mixed phase of the neutron stars are discussed. The entropy production of deconfinement phase transition can be act as a signature of phase transition, but less important and does not significantly change the thermal evolution of neutron stars. The deconfinement heat can change the thermal evolution of neutron star distinctly. ",
        "introduction": "The interior of neutron stars are considered as systems where high-density phase of strong interacting matter do exist. The chemical composition of neutron stars at densities beyond nuclear saturation densities ranging from purely nucleonic matter through hyperon or meson condensation to deconfined quark matter \\cite{ref1,ref2}. Comparing with theoretical models with the observational data, we have opportunity to constrain or understanding fundamental elements of particle and nuclear physics. Many works believed that the appearance of quark matter is implied in the structure and evolution of neutron stars \\cite{ref3,ref4,ref5,ref6,ref7,ref8,ref9}. The consequences of different phase scenarios and spin evolution for the thermal evolution of neutron stars have been reviewed in comparison with existing soft x-ray observation of thermal radiation from neutron star surfaces. They both depend sensitively on and to some extent determine the chemical composition of the stars \\cite{ref10,ref11,ref12,ref13}.  As neutron stars spin down and contract, their structure and composition change with the increasing density, and the quark matter may appear. We will discuss how the increasing density affect the spin and thermal evolution of neutron stars which contain quark matter in the core. As mentioned in the work \\cite{ref17}, the increasing density and changing chemical composition further imply additional entropy production in bulk and the release of the latent heat as particles cross any phase boundaries present. They also though that this kin of additional entropy production during deconfinement phase transition should affect the spin and thermal evolution of neutron stars.  In our former work, we discussed that the occurrence of first-order deconfinment phase transition is accompanied by the release of energy which originate from the binding energy. Therefore, we introduced the deconfinement phase transition which happened in the mixed phase of neutron star in this paper. The energy released during the deconfinement phase transition can be calculated with different equation significantly \\cite{ref10,ref11,ref12,ref15,ref16}.  In Sect. $2$ we will review the deconfinement phase transition which happens in the core of neutron stars, compare entropy production and deconfinement heat. Its application to the spin and thermal evolution of neutron stars also will be discussed. The conclusion will be given out in Sect. $3$. ",
        "conclusions": "The deconfinement phase transition during the spin-down of neutron stars have been discussed in this paper.We estimated the energy generation rate of the two different mechanisms during deconfinement phase transition. The entropy release during deconfinement phase transition which analyzed in the work~\\cite{ref17} have found do not dominate the general thermal evolution of neutron star but they do complement the standard picture. The signature of deconfinement found in this work is below the present observational sensitivity and not of sufficient strength to set apart the cooling curves with temperature versus time for hybrid stars.  The energy generated from first-order phase transitions which can act as a heating mechanism during the thermal evolution of neutron stars. It can affect the thermal evolution of compact star significantly at older age.  But the equation of state of superdense matter is still a mystery and the stars distance and ages are uncertain. Therefore, constraint on the chemical composition of compact stars purely through thermal radiation will be limited. Other observed phenomena, such as critical rotation and limited mass-radius relationship \\cite{ref20,ref21,ref22,ref23,ref24,ref25} should be used for this goal. The combination of different evolution processes, such as spin and thermal evolution of the stars which is mentioned in \\cite{ref17}, is a useful way to understand superdense matter in the compact star.  \\bigskip This work is supported by the NSFC project NO. 10773004 and the 973Program 2009CB824800."
    },
    "1005/1005.2502_arXiv.txt": {
        "abstract": "s{ We present some preliminary results from a series of extremely large, high-resolution N-body simulations of the formation of early nonlinear structures. We find that the high-z halo mass function is inconsistent with the Sheth-Tormen mass function, which tends to over-estimate the abundance of rare halos. This discrepancy is in rough agreement with previous results based on smaller simulations. We also show that the number density of minihaloes is correlated with local matter density, albeit with a significant scatter that increases with redshift, as minihaloes become increasingly rare. The average correlation is in rough agreement with a simple analytical extended Press-Schechter model, but can differ by up to factor of 2 in some regimes.} ",
        "introduction": "The properties of the first luminous objects in the universe remain a big enigma at present due to the scarcity of observational data. These objects, the first stars and galaxies, started forming very early, only about 100-200 million years after the Big Bang, and their ionizing radiation eventually completely reionized the intergalactic medium \\cite{2006PhR...433..181F}. This complex process can be studied through numerical simulations % and semi-analytical models% . The former have the advantage of being able to describe complex situations and, in particular, the non-linearities of the cosmic structures, but are expensive to run and have limited dynamic range. The latter are much cheaper to run and thus allow studies of the full parameter space, but inevitably involve many approximations and simplifications. The distribution and properties of the low-mass halos which host the first stars is an important ingredient in any semi-analytical model of reionization. Such small-scale structure potentially has very different properties from larger structures we see at later times, as they probe a very different part of the initial power spectrum of density fluctuations. Therefore, it is important to check the validity of any models and fits to the halo mass function and bias in this new regime. ",
        "conclusions": ""
    },
    "1005/1005.5322_arXiv.txt": {
        "abstract": "We study the generation of helical magnetic fields during single field inflation induced by an axial coupling of the electromagnetic field to the inflaton. During slow roll inflation, we find that such a coupling always leads to a blue spectrum with $B^2(k) \\propto k$, as long as the theory is treated perturbatively. The magnetic energy density at the end of inflation is found to be typically too small to backreact on the background dynamics of the inflaton. We also show that a short deviation from slow roll does not result in strong modifications to the shape of the spectrum. We calculate the evolution of the correlation length and the field amplitude during the inverse cascade and viscous damping of the helical magnetic field in the radiation era after inflation. We conclude that except for low scale inflation with very strong coupling, the magnetic fields generated by such an axial coupling in single field slow roll inflation with perturbative coupling to the inflaton are too weak to provide the seeds for the observed fields in galaxies and clusters. ",
        "introduction": "\\label{sec:intro}  Cosmic magnetic fields have been observed on all scales ranging from stars to near and far away galaxies and galaxy clusters \\cite{Kronberg:1993vk, Pentericci:2000mp, Battaglia:2008ex, Clarke:2000bz}. The strength of the magnetic fields observed in galaxies and clusters is typically of the order of $\\mu$Gauss. Recently, using the absence of extended GeV emission around TeV blazar gamma-rays, a lower limit of $3\\times 10^{-16}$ Gauss on the strength of intergalactic magnetic fields was derived \\cite{Neronov:2010,tavecchio-2010.2,tavecchio-2010,Dolag:2010ni}.  These observations prompt the question of the origin of cosmic magnetic fields. Have they been generated during structure formation or have primordial magnetic fields been amplified? So far, this question has no clear, satisfactory answer. Several studies~\\cite{Matarrese:2004kq, Gopal:2004ut, Takahashi:2006fea, Ichiki:2006cd, Kobayashi:2007wd, Ichiki:2007zz, Maeda:2008dv, Fenu:2010kh} of second order perturbation theory have shown that up to recombination only very weak magnetic fields of the order of $10^{-28}$ Gauss can be generated by structure formation. However, it turns out that the magnetic fields from second order perturbation theory are not strong enough to exceed the lower limit derived in Refs.~\\cite{Neronov:2010,tavecchio-2010.2,tavecchio-2010,Dolag:2010ni}. Whether such fields could have been generated later by the process of galaxy formation and then ejected into intergalactic space remains unclear.  In this work, we pursue the idea that instead magnetic fields are of primordial origin and might have been generated in the early universe. Primordial magnetic fields are interesting as they induce all three kinds of gravitational perturbations i.e. scalar, vector and tensor; all of which contribute to the Cosmic Microwave Background (CMB) temperature and polarisation anisotropies. These primordial fields also lead to non-Gaussian signals in the CMB even at the lowest order contrary to the higher order effect due to inflationary scalar perturbations. Analyses of such effects using recent cosmological data provide an upper limit of a few nGauss on the primordial magnetic fields~\\cite{Durrer:2006pc,Caprini:2009vk,Kahniashvili:2010wm,Trivedi:2010gi}. It has been argued that the electroweak or the QCD phase transitions, if they are first order, lead to the generation of magnetic fields~\\cite{Enqvist:1993kf, Joyce:1997uy, Forbes:2000gr, Boeckel:2009ej}. Also non-perturbative processes during preheating can generate significant magnetic fields with, in some cases, a helical component, see for instance \\cite{DiazGil:2008tf}. However, causality strongly constrains such fields. Their power spectrum is very blue with $B^2(k) \\propto k^2$, and therefore, their amplitude on large scales is far too small~\\cite{Durrer:2003ja, Caprini:2001nb}.  On the other hand, if the magnetic fields are produced during inflation, their power spectrum is a priori not constrained by causality but only by the specific model. Since the standard electromagnetic (EM) action is conformally invariant the fluctuations in the EM field are not amplified in the conformally flat expanding background of inflation. In order to generate magnetic fields, one needs to break conformal invariance of the EM field, e.g.~by coupling the EM field to a scalar or a pseudo-scalar field or to a curvature invariant (for an overview see, for instance, Ref.~\\cite{Subramanian:2009fu}). Typically, a term of the form $f F^2$ is considered, where $f$ is a function of time or of the inflaton and $F$ is the EM field tensor. Depending on the form of the coupling $f$, this gives rise to different magnetic field power spectra, and even scale-invariant spectra are possible (in this context, see Refs.~\\cite{Martin:2007ue, Demozzi:2009fu}). In this way, the magnetic fields can have sufficient amplitude on large scales to provide seeds for the observed fields in galaxies and clusters. However, the backreaction on the inflation dynamics, the production of gravitational waves and nucleosynthesis bounds on the amplitude of gravitational waves strongly constrain the magnetic energy density \\cite{Durrer:2006pc, Caprini:2001nb}.  Quite a different situation is encountered if a coupling to the parity-violating term $F\\Fdual$, i.e.~a term $f F\\Fdual$, is added to the standard EM action $F^2$, where $\\Fdual$ is the dual of $F$. As a consequence, magnetic helicity is generated, which is absent in the case discussed above. This has two interesting consequences: firstly, contrary to non-helical fields, helical fields evolve in the cosmic magnetohydrodynamic (MHD) plasma via inverse cascade~\\cite{Brandenburg:2004jv, Banerjee:2004df, Campanelli:2007tc}. This transfers power from small to large scales so that even blue spectra can lead to significant power on large scales. In Ref.~\\cite{Caprini:2009pr}, it was shown that the inverse cascade is not quite sufficient for helical fields generated at the electroweak phase transition~\\cite{Vachaspati:2001nb}, but it might work for magnetic fields from inflation. Secondly, helical magnetic fields leave a very distinct signature as they violate parity symmetry. This leads to observable effects, e.g.~correlations between the anisotropies in the temperature and B-polarisation or in the E- and the B-polarisations in the CMB~\\cite{Caprini:2003vc}. Furthermore, they induce helical gravitational waves~\\cite{Caprini:2003vc} which might be observable~\\cite{Seto:2008sr}.  Some consequences of primordial helical magnetic fields and their generation from primordial helicity have been studied in the past~\\cite{Cornwall:1997ms, Field:1998hi}. The interactions of helical magnetic fields and axions have also been investigated~\\cite{Campanelli:2005jw, Campanelli:2005ye}. Recently, the generation of helical magnetic fields during inflation in specific models has been studied, e.g.~in Ref.~\\cite{Anber:2006xt}, helical magnetic fields from N-flation were investigated. In this case, the large number of pseudo-scalar fields driving inflation effectively leads to a large coupling, $f\\propto\\sqrt{N}$, to the $F\\Fdual$ term. In Ref.~\\cite{Campanelli:2008kh}, some toy models for the coupling were analysed where $f$ was taken to be a power law function of $k\\tconf$ ($k$ being the wave number and $\\tconf$ the conformal time).  In this paper, we study magnetic fields generated by an axial coupling of the form $f(\\phi)F\\Fdual$ during inflation, where $\\phi$ is the inflaton. We consider two different forms of the coupling function and show that, contrary to a non-helical coupling of the form $f(\\phi)F^2$, a helical coupling always leads to a spectral index $n=1$ for $B^2(k) \\propto k^n$, as long as slow roll inflation is considered. We derive the condition for  the theory to be perturbative, i.e. the free part of the action dominates over the interaction. Of course this is not necessarily true, but at least naively if not, we can no longer trust our calculation of particle creation out of the quantum vacuum which is based on perturbative quantum field theory. We estimate the magnetic energy density as a function of scale and show that backreaction is typically small. These conclusions are valid for any reasonable coupling function $f(\\phi)$. We confirm our analytical results numerically. Furthermore, we study the effects of a short deviation from slow roll on the magnetic field spectrum and show that such deviations, if kept within the bounds permitted by the CMB data, do not strongly modify its shape. Even though the inverse cascade in the radiation dominated era after inflation does move power to larger scales, the final strength of the magnetic field on cosmologically interesting scales is still insufficient to provide seeds for the observed magnetic fields in galaxies and clusters, except if the inflation scale is low, $T_\\ast  < 10^4$GeV and the axial coupling is very strong. Even if we assume very efficient dynamo amplification.  The remainder of this paper is organised as follows. In the next section, we introduce the axial coupling of the EM field to the inflaton, derive the field equations, discuss the background evolution and the slow roll approximation and derive the linear perturbation equations of the inflaton and the EM vector potential. In~\\secref{sec:modeeq}, we discuss the evolution equation of EM quantum fluctuations during inflation and compare the helical to the normal case. In~\\secref{sec:slowroll}, we derive a condition on the chiral coupling by requiring the theory to be perturbative, investigate different coupling functions in the slow roll approximation, and solve the evolution of the EM fluctuations analytically. In~\\secref{sec:numerical}, we discuss the consequences of a brief violation of the slow roll approximation on the magnetic field spectrum. In~\\secref{sec:interpretation}, we study the evolution of the magnetic field during the radiation era after inflation and determine the final spectrum after the inverse cascade. Finally, in~\\secref{sec:conclusions}, we conclude with a few comments on our results. Three appendices contain some details on the quantisation of the vector potential, the condition on the coupling function from the perturbativity of the theory and the asymptotic behaviour of the Coulomb wave functions, respectively.   \\paragraph{Notation and units:} We work in a metric with signature \\mbox{($-$ + + +)}. For tensor components, Greek indices take values $0\\ldots3$, while Latin indices run from $1$ to $3$. The components of spatial $3$-vectors with respect to a comoving basis are denoted in bold face. We employ Heaviside-Lorentz units such that $c=\\hbar=k_B=\\epsilon_0=\\mu_0=1$. The reduced Planck mass is defined as $m_P=(8\\pi G)^{-1/2}$. We normalise the cosmic scale factor to unity today so that the comoving scales become physical scales today. ",
        "conclusions": "\\label{sec:conclusions}  In this work, we have studied the generation of helical magnetic fields during inflation by adding a parity-violating term of the form $f(\\phi)F\\Fdual$ to the action. For two specific choices of the coupling function $f(\\phi)$, namely a power law and an exponential form, we found that, during slow roll inflation, the power spectrum of the generated magnetic fields is always blue with spectral index $n=1$. The aspect  responsible for this result is that the helical coupling term in the evolution equation for the Fourier modes of the vector potential is active only around horizon crossing of a given mode and it is not varying rapidly during slow roll inflation. As a consequence, in slow roll inflation all Fourier modes are amplified in the same way and the final shape of the power spectrum is unchanged with respect to the initial one.  We have also studied the effects of a short deviation from slow roll. If kept within the bounds allowed by the CMB data, such a departure from slow roll does not strongly modify the overall shape of the magnetic field power spectrum. In principle, a very steep coupling function, e.g.~a double-exponential $f(\\phi)=f_0\\exp[\\exp(\\al\\phi/m_P)]$, can lead to a different spectrum, $n\\neq 1$. However, if we want to satisfy perturbativity, $|\\fN|< 1$, for such a behaviour to be valid throughout inflation, a ridiculously small pre-factor $f_0$ is required. In the absence of a convincing physical motivation, we have not investigated such an extreme case.  After inflation and reheating, an inverse cascade sets in and moves power from the correlation scale $k_*\\simeq \\H_*$ to larger scales, which is schematically represented in \\figref{f:casc2}. We find that typical values for the reheating temperature, e.g.~$T_*\\simeq 10^{14}\\,\\GeV$, can only lead to magnetic fields of the order of $10^{-40}\\,\\Gauss$ on scales of 0.1 Mpc and smaller values are obtained for smaller reheating temperatures.  {Unless $\\CS(f_N)\\gg 1$ and the reheating temperature is sufficiently low, $T_* \\le 10^4$GeV, }  these field amplitudes are largely insufficient for dynamo amplification~\\cite{Brandenburg:2004jv}.  This result is quite generic. A power law or an exponential coupling to the $F\\Fdual$ term during slow roll inflation will generate sufficiently strong seeds for the large scale magnetic fields in galaxies and clusters without also leading to backreaction only if the inflation scale is sufficiently low, of the order of the electroweak scale.  Even though for helical fields an inverse cascade does move power to larger scales, this is still insufficient in most cases. For typical inflation scales, to obtain sufficient fields on large scales after the inverse cascade, they must have been so large on small scales after inflation that backreaction cannot be neglected. Another possibility might be a very steep coupling which can lead to a different, $n\\neq 1$ spectrum. However, only if the spectrum is close to scale invariant, it can have significant amplitudes on large scales without significant backreaction from small scales.  Scenarios of inflationary magnetogenesis are often constrained by requiring that the backreaction of the generated magnetic field on the background evolution is small. Since the perturbations in the scalar field are affected by the presence of a non-minimal coupling as indicated in \\eqref{e:dephi}, it will be interesting to study the backreaction effects of the EM perturbations on the evolution of the primordial curvature perturbations. This may provide another tool to constrain inflationary scenarios of magnetogenesis. This issue will be addressed in a future project~\\cite{WeInPrep}."
    },
    "1005/1005.0380_arXiv.txt": {
        "abstract": "The XENON100 experiment, in operation at the Laboratori Nazionali del Gran Sasso in Italy, is designed to search for dark matter WIMPs scattering off $62\\1{kg}$ of liquid xenon in an ultra-low background dual-phase time projection chamber. In this letter, we present first dark matter results from the analysis of 11.17 live days of non-blind data, acquired in October and November 2009. In the selected fiducial target of $40\\1{kg}$, and within the pre-defined signal region, we observe no events and hence exclude spin-independent WIMP-nucleon elastic scattering cross-sections above $3.4\\times10^{-44}\\1{cm}^2$ for $55\\1{GeV/c^2}$ WIMPs at $90\\percent$ confidence level. Below $20\\1{GeV/c^2}$, this result constrains the interpretation of the CoGeNT and DAMA signals as being due to spin-independent, elastic, light mass WIMP interactions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.2982_arXiv.txt": {
        "abstract": "Magnetic fields appear to be a generic feature of the early universe and are a natural source of secondary CMB non-Gaussianity. In recent years the statistical nature of the stresses of a primordial magnetic field has been well studied. In this paper we confirm and extend these studies at one- and two-point level, and present analytical results for a wide range of power-law spectra. We also consider two non-power law cases of interest: a blue spectrum with an extended damping tail on small scales, which could be generated by the non-linear mixing of density and vorticity; and a red spectrum with a damping tail on large scales. We then briefly consider the CMB impacts that result from such fields. While this paper focuses on the one- and two-point moments, the techniques we employ are designed to ease the analysis of the full bispectra induced by primordial magnetic fields. ",
        "introduction": "Large-scale magnetic fields appear to be inevitable. Fields have been observed on galactic and cluster scales and likely exist on supercluster scales \\citep{Kronberg:1993vk,Grasso:2000wj,Giovannini:2003yn,Xu:2005rb}. There are a variety of potential generation mechanisms in both the early and late universe, many of which may well have been in operation, ranging from so-called Biermann batteries (e.g. \\citep{Biermann:2003xh,Hanayama:2005hd}) to fields generated at reionisation \\citep{Subramanian:1994,Gnedin:2000ax,Langer:2002tx,Langer:2005gw}. The number of well-motivated magnetogenesis mechanisms would appear to make the existence of magnetic fields across a large range of scales inevitable. In particular, plasma processes in the early, pre-recombination universe \\citep{Gopal:2004ut,Matarrese:2004kq,Takahashi:2005nd,Ichiki:2006cd,Siegel:2006px,Kobayashi:2007wd,Maeda:2008dv} wherein linear scalar perturbations drive nonlinear vorticity, producing currents in the proton/electron plasmas, would suggest that at least some magnetic fields before recombination are inevitable even if the detailed predictions in the various models differ somewhat. The fields extent in the present universe need not necessarily be primordial; the fields observed on cluster scales can be accounted for in principle by non-linear processes before recombination, by production at reionisation or at recombination \\citep{Hogan:2000gv,Berezhiani:2003ik}, and it seems probable that each of these processes have contributed. There also many exist potential mechanisms in the extremely early universe, such as magnetogenesis at a phase transition (for example \\citet{Baym:1995fk,Martin:1995su,Hindmarsh:1997tj,Boyanovsky:2002wa,Kahniashvili:2009qi}).  However, all of these fields will have a distinctive, blue power spectrum \\citep{Durrer:2003ja}, and the most recent bounds on fields present at or before recombination rather favour a red spectrum \\citep{Yamazaki:2010nf,Paoletti:2010rx}. It seems reasonable, therefore, to also consider mechanisms that could produce such red fields. To achieve a near-scale invariant spectral index these will necessarily be produced during an inflationary epoch \\citep{Durrer:2003ja}, and there is a large variety of potential genesis mechanisms that generate a field from the breaking of the conformal invariance of the electromagnetic field, with \\citep{Turner:1987bw,Bassett:2000aw,Prokopec:2004au,Giovannini:2007rh,Campanelli:2007cg,Bamba:2008my} being a few examples.  Constraints on a field produced during the inflationary epoch have naturally focused on the cosmic microwave background (CMB) (\\citet{Barrow:1997mj,Clarkson:2002dd,Giovannini:2009fu,Yamazaki:2010nf,Paoletti:2010rx}, for example) and have typically found that the field strength today is at most of the order of nanogauss. A stochastic background of primordial magnetic fields, such as would be reasonably expected from an inflationary process, impacts on the CMB in two chief ways. Firstly the magnetic field itself acts as an additional radiative fluid which, should there be enough energy in the field, impacts on the expansion rate of the universe. However, by far the most significant impact comes from the interactions of the magnetic field and the ionised plasmas extent before recombination, and the scattering between the magnetic fields and the geometry. There have been numerous studies into magnetised cosmological perturbation theory; for some interesting studies see for example \\citep{Jedamzik:1996wp,Subramanian:1997gi,Tsagas:1999ft,Mack:2001gc,Giovannini:2003yn,Tsagas:2004kv,Banerjee:2004df,Giovannini:2005jw,Barrow:2006ch,Brown:2006wv,Giovannini:2007qn,Kojima:2009ms} and their references. Clearly the presence of the magnetic field will impact on the CMB angular power spectra. While the impact on the temperature auto-correlation and on the temperature/polarisation cross-correlation will be small except at larger multipoles, a magnetic field also excites $B$ mode polarisation. The $B$ mode angular power spectrum would therefore be of great interest for the study of primordial magnetic fields. See for example \\citep{Subramanian:1998fn,Durrer:1998ya,Koh:2000qw,Kahniashvili:2000vm,Mack:2001gc,Lewis:2004ef,Kahniashvili:2006hy,Kahniashvili:2008hx,Yamazaki:2008gr,Finelli:2008xh,Paoletti:2008ck,Bonvin:2010nr} and their references for studies of the magnetised CMB angular power spectra.  The $B$ mode spectrum produced by a primordial magnetic field is, however, likely to be nearly degenerate with the gravitational wave spectrum produced by inflationary models \\citep{Lewis:2004ef} and even with forthcoming results from for example the Planck and QUIET experiments it is possible that we will be unable to disentangle the impact of magnetic fields from other contributions and from foreground signals. An alternative is provided by the non-Gaussianity of the magnetic field, studied in for example \\citep{Brown:2005kr,Brown:2006wv,Seshadri:2009sy,Caprini:2009vk}; since the stress-energy tensor and the Lorentz forces generated by a magnetic field are non-linear, the statistics imparted on the perturbations by the magnetic field are necessarily non-Gaussian, regardless of the statistics of the underlying magnetic field. Furthermore, we would expect the greatest impact to be on large scales, observations of which with the WMAP satellite are already cosmic-variance limited \\citep{Komatsu:2010fb}. Study of the magnetised CMB bispectrum is, however, unfortunately a difficult issue. It requires knowledge of the full 3-point moment of the stress tensor of the magnetic field (the ``shape function'' or ``intrinsic bispectrum'') and to date only a few particular configurations have been studied. While these cases are probably sufficient for nearly scale-invariant magnetic fields, for which power appears to be concentrated on the so-called local configuration (with the bispectrum geometry set by three wavevectors $\\mathbf{k}_1\\approx -\\mathbf{k}_2$ and $\\mathbf{k}_3\\approx 0$), it has not been possible to fully test this. Furthermore, for the blue cases expected to arise in non-inflationary magnetogenesis models neither the local, nor indeed any other, configuration will necessarily dominate and knowledge of the full intrinsic bispectrum will be required.  To consider the impact of a magnetic field on the CMB, then, we require a clear understanding of the statistical nature of the stress tensor (and, equivalently, the Lorentz forces). In \\citep{Brown:2005kr} and \\citep{Brown:2006wv} (hereafter BC05 and B06) we studied the 1-, 2- and a special case of the 3-point moments of the stress tensor of a Gaussian magnetic field with a power law power spectrum with index $n_B$, employing a mixture of realisations of the magnetic fields and numerical integration. The nature of the 2-point moments changes depending on whether $n_B>-3/2$ or $n_B<-3/2$ and we considered two cases in particular. For a field with a flat spectrum we found the 2-point moments with both methods, demonstrating good agreement and describing the gross features of such ``ultra-violet'' fields. For a field with a strongly red spectrum we only employed realisations and could only draw the broadest of conclusions. In B06 we further considered a ``damped causal'' power spectrum, resembling that produced by plasma processes in the pre-recombination universe, and the statistics of the Lorentz forces (as subsequently studied in, for example, \\citep{Paoletti:2008ck}). More recent studies have improved the approach to numerical integration \\citep{Yamazaki:2008gr}, and extended it to a wider range of spectral indices; and have even found analytical solutions for particular cases \\citep{Finelli:2008xh,Paoletti:2008ck}. We argue in this paper that the most important spectral indices are at $n_B=2$, $n_B=0$, $n_B=-3/2$ and $n_B=-5/2$, all of which are included in their solutions.  However, while this means that we now have solutions from analysis, numerical integration and from statistical realisations work remains to be done. In principle the realisations are the most general approach, and can be applied to any field, possessing any statistical nature and any power spectrum. However, they are severely compromised by the mode coverage on large scales associated with the finite size of the grid. The extent of this issue was not considered before. Due to its reliance on Wick's theorem, numerical integration can at present only be applied to fields with Gaussian statistics, although arbitrary power spectra can be employed. Our approach in BC05 and B06 was restricted to blue power-law spectra with $n_B>-3/2$ which is overly restrictive. The approach of \\citep{Yamazaki:2008gr} is more general but relatively unwieldy and has been applied only to power-law spectra. Finally, while \\citep{Paoletti:2008ck} found analytic solutions for spectral indices of interest, they are missing in particular the case $n_B=-2$. As there are only two analytic solutions in the region $n_B\\in(-3,-3/2)$, at $n_B=-5/2$ and $n_B=-2$, this is an important omission. The subsequent approach to CMB observables, given these spectra, has been rather piecemeal and model-dependant.  It is our intention in this paper to reconsider many of these issues, always focussing on retaining as much generality and flexibility as possible. In particular, we wish to explicitly decouple the consideration of the the intrinsic magnetic statistics from the evaluation of the CMB angular power spectra; in this manner, one may build the statistical nature of the magnetic stress tensor in any manner one wishes, and wrap it onto the CMB employing independently-derived analytical or numerical transfer functions.  Further, we wish to approach the issue of the intrinsic statistics in an extensible manner, unifying to some degree the approach to analysis and to numerical integration. In doing so we extend the set of useful analytic solutions, finding solutions for all $n_B$ between $n_B=3$ and $n_B=-5/2$. While we will not generalise beyond underlying Gaussian statistics for the magnetic field, we also consider two non-power law power spectra of interest. The first is a damped causal spectrum similar to that considered in B06, and the second resembles a strongly red field produced during an inflationary epoch with an extended decaying tail on large scales. We find analytical solutions for the former and rely on numerical integration for the latter. We then consider the CMB, employing a clearly-defined approximation based on the generic form of a magnetised transfer function. In this paper we consider for brevity only the 1- and 2-point moments but we emphasise throughout the paper that we have constructed our techniques to be immediately applicable to the study of magnetic bispectra, to which we will turn in a forthcoming study.  We start in section \\ref{MagneticFields} by introducing our toy model of primordial magnetic fields and briefly discussing the normalisation of their power spectrum. In section \\ref{OnePoint} we briefly consider the one-point statistics of the stress tensor before considering in somewhat more detail the two-point moments in section \\ref{2Points}. In section \\ref{CMB-General} we present our general approach to the magnetised CMB angular power spectra, and in section \\ref{CMB-Results} apply this to the CMB temperature auto-correlation generated by tensor perturbations. We briefly conclude in section \\ref{Conclusions}. Appendix \\ref{Appendix-SpectraSolutions} presents those analytical solutions not previously presented in \\cite{Paoletti:2008ck}, appendix \\ref{Appendix-LargeScale} presents the forms on the largest scales, and appendix \\ref{Appendix-Yamazaki} briefly addresses a comment in \\citep{Yamazaki:2008gr} concerning our previous work. Latin indices refer to spatial dimensions and are raised and lowered with the Minkowski metric $\\eta_{ij}=\\mathrm{diag}(1,1,1)$. ",
        "conclusions": "\\label{Conclusions} In this paper we have considered in detail the statistics of the stress tensor of a tangled primordial magnetic field at the one- and two-point levels, employing a mixture of analytical techniques, numerical integration, and statistical realisations. Assuming Gaussian statistics and a power-law spectrum for the magnetic field, we are restricted to the region $n_B>-3$; a field with $n_B=-3$ would be exactly scale-invariant. Fields with $n_B\\geq 2$ can be produced by causal processes. We have examined the one-point moments, extending the results of BC05 and B06 to a wide range of spectral indices and a much higher grid-resolution.  We have verified that the isotropic pressure of the magnetic field is a $\\chi^2$ field to a very high degree, as has been argued before. We also considered the anisotropic pressure, which runs with spectral index. In particular, for indices $n_B\\geq -3/2$ the statistics of the anisotropic pressure are constant, with a skewness $\\gamma_1\\approx -0.24$ and a kurtosis $\\gamma_2\\approx 1.1$. However, in the region $n_B<-3/2$ both the skewness and the kurtosis increase as $n_B\\rightarrow -3$. There is a strong grid dependence for such steeply-tilted spectra which requires us to operate on grids of side-length $l_\\mathrm{dim}=512$, which arises from the power being piled onto a small number of modes on the finite grid. This infra-red divergence is inevitable -- and re-emerges in the study of the 2-point moments -- but we can nevertheless predict that $\\gamma_1\\lesssim 1$ and $\\gamma_2\\lesssim 2$ at $n_B=-3$. Perhaps more importantly, we have demonstrated that at the one-point level the field with $n_B=-3/2$ separates the parameter space into two regions in which the anisotropic stress behaves quite differently.  Turning to the two-point moments, our chief focus has been to construct a formalism which can be equally applied to both analytic and to numeric evaluation of the stress spectra -- and which, more importantly, will be readily adaptable to the calculation of magnetic bispectra. We then applied this to the two-point moments of a power-law magnetic field, finding analytic solutions existing for all integer and half-integer $n_B\\in[-5/2,3]$. While some of these solutions were previously found in \\citep{Paoletti:2008ck}, others were not. As has been known for some time, the parameter space is split in two, again at $n_B=-3/2$. For $n_B>-3/2$ the integrals are readily solved, and across the range of $k$, the stress spectra are technically ``decoherent'', in that the amplitude and the gradient at each $k$ depend explicitly on the damping scale, which is time-dependent. On the very largest scales, for $k\\lesssim k_\\mathrm{Coh}\\approx k_c/100$ the cross-correlation of the isotropic and anisotropic pressures vanishes (and is an anti-correlation on intermediate scales, except at $n_B=-1$), while the other correlations tend towards white noise. The ratio between the various correlations is constant across $n_B$, and is equivalent to the ``causal'' ratios found in \\citep{Turok:1997gj}. Conversely, for $n_B<-3/2$, there is a large region in which the statistics are ``coherent'', for wavenumbers up to $k_\\mathrm{Coh}\\approx k_c$. On these scales the power spectra scale as $\\mathcal{P}(k)\\propto k^{2n_B+3}$ as has been appreciated for some time; it is worth noting that it is only for integer and half-integer $n_B$ that this results in integer powers of $k$. Interestingly, the cross-correlation between the scalar pressures is of the same order of magnitude as the auto-correlations. Above $k_\\mathrm{Coh}$ the statistics remain entirely decoherent. In particular, we presented a solution for the case with $n_B=-2$. Since we cannot find a solution at $n_B=-3$, there are only two solutions in this ``coherent'' or ``infra-red'' region available, $n_B=-5/2$ and $n_B=-2$. Since these values remain in the centre of the allowed constraints, it is important that we possess the solutions. Finally, the field with $n_B=-3/2$ behaves quite differently. In this case the cross-correlation of the scalar pressures is a positive correlation, and white noise on large scales, while the auto-correlations all diverge logarithmically with $k$. Numerical integrations employing Monte-Carlo techniques confirmed the analytic solutions to an extremely good accuracy.  From our analysis, we then not only have an approach to analytical and numerical integrations that can be extended with ease to bispectra, but we have also clearly identified the important regions of the spectra solutions, broadly governed by a coherence scale $k_\\mathrm{Coh}$.  We then considered two non-power law magnetic power spectra -- the first, a ``damped causal'' field has a spectral index $n_B=2$ on large scales but an extended damping tail on smaller scales, and resembles a field produced by a causal mechanism such as the nonlinear vorticity of the electron/proton plasmas; while the second, an ``IR-controlled'' field, has a spectral index $n_B=-5/2$ on relatively small scales but decays on large scales, and while it could in principle be produced by a carefully-chosen inflationary scenario is more concretely employed to model the infra-red cut-off of a sample realisation. Analytic solutions exist for the damped causal field we chose. The damped causal stress spectra resemble the standard causal case with the gross features moved to a smaller wavenumber corresponding to the peak of the magnetic power spectrum. The IR controlled spectra, in contrast, tend towards white-noise for $k<k_S$ where $k_S$ is related to the scale at which large-scale magnetic power begins to decay, but closely resemble the standard infra-red case on larger scales, albeit with a suppression of power for all $k<k_\\mathrm{Coh}$. For both of these scenarios, though, the cross-correlation significantly differs from the undamped cases, which might in principle allow us to distinguish between them.  We also considered the power-law spectra employing statistical realisations, running on grids of side-length $l_\\mathrm{dim}=256$. While the results in the ultra-violet regime with $n_B>-3/2$ are in extremely good agreement with the analytical solutions, the limitations of the approach are quickly apparent: we are limited to a minimum wavenumber $k/k_c\\approx 4/l_\\mathrm{dim}\\approx 10^{-3}$-$10^{-2}$. While this is close to the white-noise regime, one would not necessarily wish to na\\\"ively rely on the output at such a large scale, particularly as the error at the minimum wavenumber is naturally rather large. When considering $n_B<-3/2$ the limitations of the realisations are even clearer; there is an extremely strong impact from the infra-red damping. While working on increasingly large grids eases the problem this rapidly becomes unwieldy and unrealistic. While realisations remain in principle the most flexible approach to the statistics of early universe magnetic fields, not least as they can easily take into account magnetic fields with non-Gaussian statistics, they should be treated with caution.  Finally, we outlined an approach to the CMB which will be of use in the study of the CMB angular bispectrum arising from magnetic perturbations. Working with the tensor perturbations sourced by a magnetic field we considered the temperature auto-correlation. Using that the transfer function is steeply tilted to large scales, we argued that generally only the regions $k<k_\\mathrm{Coh}$ contribute to the CMB integrations. The stress spectrum can then with fair generality be written as a simple power law, where the amplitude $P_\\star$ at a convenient pivot scale $k_\\star$ can be recovered from an analytic solution, from a numerical integration, or from a statistical realisation. Assuming that the approximate form for the tensor transfer function is accurate, exact general solutions for the CMB signals can then be found, and specialised to particular power spectra with a suitable choice of $P_\\star$ and $\\alpha$. Where this is not true, an approach similar to that which we took for the IR-controlled field can be applied: the stress spectrum can be split into two or more regions, each characterised by a different $\\alpha_i$. The solutions that result are unwieldy combinations of hypergeometric functions, but in most cases are likely to reduce to simpler forms for realistic parameters.  In addition to reproducing the previously-known forms for the CMB angular power spectrum from magnetised tensor modes and power-law magnetic fields, we considered the damped causal and the IR-controlled fields. For the parameters we chose in this paper, which seem reasonable, the signals turn out to be indistinguishable from those from the standard causal and inflationary cases respectively. However, we have presented the expressions for a wider variety of parameters. We have also evaluated the amplitudes of such fields when smoothed on a comoving scale $k_\\lambda$.  This approach to the CMB entirely separates the form of the stress spectrum from the transfer functions in a transparent manner. Furthermore, the arguments we presented do not rely on the specific analytic forms of the stress spectra; rather, they rely on the form of the magnetised transfer functions. Naturally, these do not differ between the angular power spectra and the angular bispectra. The bispectrum can be described as with the coordinate system $\\{k,r,\\phi\\}$ and foliated into planes of $\\{k,r\\}$; in each of the 1D lines of constant-$r$ through these planes it appears \\citep{Brown:2005kr,Brown:2006wv,Caprini:2009vk,Seshadri:2009sy} that the bispectra on large scales behave as either white-noise in $k$ or with a power law in $k$. Our approach to the CMB angular power spectra, therefore, can equally be applied to the magnetised CMB angular bispectra, which are otherwise extremely challenging to consider. This will be the focus of future study."
    },
    "1005/1005.5100_arXiv.txt": {
        "abstract": "In this paper, we present a model for the effects of the tachocline on the differential rotation in the solar convection zone.  The mathematical technique relies on the assumption that entropy is nearly constant (``well-mixed'') in isorotation surfaces both outside and within the tachocline. The resulting solutions exhibit nontrivial features that strikingly resemble the true tachocline isorotation contours in unexpected detail. This strengthens the mathematical premises of the theory. The observed rotation pattern in the tachocline shows strong quadupolar structure, an important feature that is explicitly used in constructing our solutions. The tachocline is treated locally as an interior boundary layer of small but finite thickness, and an explicit global solution is then constructed.  A dynamical link can thus be established between the internal jump in the angular velocity at the tachocline and the spread of angular velocities observed near the solar surface. In general, our results suggest that the bulk of the solar convection zone is in thermal wind balance, and that simple quadrupolar stresses, local in radius, mediate the tachocline transition from differential rotation to uniform rotation in the radiative interior. ",
        "introduction": "Recent work suggests that the isorotation contours in the bulk of the solar convection zone (hereafter SCZ), correspond to the characteristic curves of the vorticity equation in its ``thermal wind'' form (Balbus 2009, hereafter B09; Balbus et al. 2009, hereafter BBLW; Balbus \\& Weiss 2010). In its undeveloped form, this equation relates entropy gradients to angular velocity gradients. The mathematical existence of the characteristics demands, therefore, some sort of functional relationship between the entropy $S$ and angular velocity $\\Omega$. The theory presented in the above papers is predicated upon the notion that within constant $\\Omega$ surfaces, the dominant, long-lived convective cells carry out their task of mixing entropy with great efficiency. This means that within a constant $\\Omega$ surface, the excess entropy gradient above and beyond the minimal radial gradient needed to maintain the convective state is, in essence, eliminated.   The remaining ``residual'' entropy is thus uniform in constant $\\Omega$ surfaces.  More explicitly, the residual entropy $S'\\equiv S-S_r$ is formed by subtracting a function of spherical radius $r$ (something close to the driving background entropy that has been averaged over angles), and $S'$ is set equal to a function of $\\Omega$.  This is a powerful constraint, providing the necessary mathematical link for interpreting the solar isorotation contours as trajectory characteristics of a quasilinear partial differential equation. BBLW advances several different lines of argument (heuristic, numeric, and goodness-of-fit) supportive of this picture.  Originally an MHD basis for the $S-\\Omega$ connection was put forth (B09), but the greater simplicity and generality, as well as the broad agreement with many different lines of argument, together suggest primacy for the BBLW mechanism.  An attractive feature of the theory presented in BBLW is that it provides simple and intuitive answers to not-so-trivial questions.  For example, why can't the solar convection zone rotate on cylinders?  Because if it did, then $\\Omega$ would necessarily depend upon the spherical colatitude angle $\\theta$, therefore the entropy would also depend upon $\\theta$, and this is inconsistent with a vanishing $z$ gradient in the thermal wind equation (TWE; cf. equation [2] below). Why then is $\\Omega$ dominated by its $\\theta$ gradient?   Now we argue the other way: because the residual entropy has its dominant radial gradient removed (physically by mixing, mathematically by a simple subtraction), the relative importance of its $\\theta$ gradient has been enhanced.  But residual entropy shares the same isosurfaces as $\\Omega$.  In numerical simulations, the fact that imposed latitudinal entropy gradients enforced at an interior boundary lead also to latitudinal gradients in $\\Omega$ (Miesch, Brun, \\& Toomre 2006) becomes tautological: latitudinal $S$-gradients lead inevitably to latitudinal $\\Omega$-gradients because these gradients are always counteraligned.  In its current formulation, BBLW theory addresses the differential rotation in the Sun away from the tachocline.  This is nothing if not prudent: the dynamics of the solar tachocline is a notoriously controversial and difficult topic (e.g. Hughes, Rosner, \\& Weiss 2007). Yet, in studying the isorotation contours within the tachocline, one is struck both by the highly localized nature of the distortion, and by the relative simplicity and uniformity of the their appearance.  In contrast to the bulk of the convection zone and the outer boundary layer, there is a sense of inevitability to these curves. Perhaps their gross form might not be sensitive to the complex details of the internal dynamical structure of the tachocline, any more than the gross adiabatic temperature profile of the SCZ depends upon the details of the convective heat flux.  We might look instead to the exigencies of local forcing and to the strikingly quadrupolar structure of the region. The resulting mathematical demands of joining the isorotation contours in such a forced region to the overlying convective zone contours treated in BBLW may be sufficiently restrictive that the role of turbulence is tightly confined.  In its current form, BBLW theory has too much ``free play:''  it needs to be restrained. In this work, we will develop this point of view, and pursue its consequences. ",
        "conclusions": ""
    },
    "1005/1005.1453_arXiv.txt": {
        "abstract": "X-ray observations of  galaxy clusters at high redshift ($z\\gtrsim 0.5$) indicate that they are more morphologically complex and less virialized than those at low-redshift.  We present the first subarcmin resolution at 18 GHz observations of the Sunyaev-Zel'dovich (SZ) effect for Cl~J0152-1357 using the Australia Telescope Compact Array. Cl~J0152-1357 is a massive cluster at redshift $z=0.83$ and has a complex structure including several merging subclumps which have been studied at optical, X-ray, and radio wavelengths. Our high-resolution observations indicate a clear displacement of the maximum SZ effect from the peak of X-ray emission for the most massive sub-clump. This result shows that the cluster gas within the cluster substructures is not virialised in Cl~J0152-1357 and we suggest that it is still recovering from a recent merger event. A similar offset of the SZ effect has been recently seen in the `bullet cluster' by Malu et al.  This non-equilibrium situation implies that  high resolution observations are necessary to investigate galaxy cluster evolution, and to extract cosmological constraints from a comparison of the SZ effect and X-ray signals. ",
        "introduction": "The Sunyaev-Zel'dovich (SZ) effect is the spectral distortion of the Cosmic Microwave Background (CMB) due to inverse Compton scattering of CMB photons off a cloud of electrons (Birkinshaw 1999). The increase of the photon energy implies a decrease of the CMB brightness at frequencies below 218 GHz and an increase at higher frequencies. The amplitude of the distortion depends only on the properties of the electron cloud: in the case of thermal electrons it is proportional to the integral of the electron pressure (i.e. $\\Delta T_{SZ}\\propto n_eT_e$) along the line of sight, and is independent of distance. Galaxy clusters are massive structures permeated by hot dense ionized gas that preserves the same baryonic fraction of the epoch of the cluster virialization, so they are the ideal targets for SZ effect observations and to infer pieces of information about cosmology.  The X-ray luminosity from the hot cluster gas has a different dependence on gas properties (i.e. $L_X\\propto n_e^2T_e^{1/2}$) and it is dependent on the cluster redshift. For virialized clusters, where the distribution of density and temperature are easy to model with regular forms (e.g. isothermal and spherically symmetric) the different distance dependence, provides a powerful method to extract redshift and cosmological parameters by combining the two signals (Cooray 1999, Carlstrom, Holder \\& Reese 2002, Molnar, Birkinshaw \\& Mushotzky 2002).  The comparison of the two signals is also a particularly effective tool to estimate the gas mass but the different dependence on density makes X-ray more susceptible to the gas clumping factor $C=\\langle n_e^2\\rangle^{1/2}/\\langle n_e\\rangle$, that may be $\\gg 1$ at high redshift. The SZ effect signal also is a better tracer of less dense hotter region.  Current theories of structure formation predict that clusters form hierarchically via merger of smaller structures (Borgani \\& Guzzo 2001). X-ray observations (Jeltema et al. 2005, Maughan et al. 2008) of high redshift ($z\\gtrsim 0.5$) clusters revealed that they are more morphologically complex, less virialized and dynamically more active than low-redshift clusters. Studies of $z>0.8$ clusters show clumpy and elongated structures that suggest that they are close to the epoch of cluster formation (Rosati et al. 2004).  The study of high redshift cluster SZ effect can provide constraints on the theories of cluster formation and evolution. However, for these objects the temperature and density distributions might be so complex that the comparison of X-ray and SZ effect signals is seriously compromised. If this is the case the information about cosmology or cluster evolution will be misleading.  In this Letter we present the first subarcminute resolution 18 GHz observations of SZ effect with the Australia Telescope Compact Array using the new wide bandwidth 4 GHz correlator (Ferris \\& Wilson 2002) for a high redshift cluster. The properties of the cluster and a summary of its existing observations in radio, optical and X-ray are described in \\S \\ref{sec:cl0152}. In \\S \\ref{sec:atdata} we present our new observations. Their comparison with other bands and the implications of our results are discussed in \\S \\ref{sec:discussion}. To estimate distances we have adopted a flat $\\Lambda$CDM cosmology with $\\Lambda=0.73$ and ${\\rm H}_0 = 71\\, {\\rm km\\, s^{-1}\\, Mpc^{-1}}$. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{The comparison with other bands}  The angular resolution of both the X-ray and radio images is sufficient to resolve the offset structures so the displacement can not be a resolution effect. We also know from the agreement in radio and X-ray positions of the two AGN just outside the cluster that the two coordinate systems are aligned to at least arcsecond accuracy.  However, from Fig. \\ref{fig:comp_x} it is clear that the SZ effect  from the NE subclump is shifted $\\sim45$ arcsec ($\\sim$342 kpc) towards the North with respect to the X-ray peak of emission and is about 30 arcsec West of the main galaxy oncentration (see Fig. \\ref{fig:comp_gal}). A similar offset of the SZ effect has been recently seen in the `bullet cluster' by Malu et al. (ApJL submitted).  The cluster is very dynamically complex, and the northern subcluster is at the intersection between two merger axes and may have undergone mergers along both axes recently  The position of the maximum SZ effect ($\\Delta T \\propto  n_e T_e$) should trace the position of the gas pressure peak of this sub-clump but the X-ray (free-free) emission ($L_X \\propto n_e^2 T_e^{1/2}$),  which also agrees with the weak-lensing-based mass distribution (Jee et al. 2005), is displaced 45 arcsec to the South-East so we need to explain this difference.  It is possible that the quality of X-ray data in the region of the SZ peak is relatively low, so that the temperature ($\\sim4.4\\pm1$ keV in Maughan et al. 2006) appears a bit lower than the peak of X-ray signal, triggering the displacement. Deeper X-ray observations will be considered to discard this possibility, but we can assume that at a first approximation we are observing a genuine displacement of the two signals due to the gas property distributions.  The group of galaxies at $z \\sim 0.6$ could affect the X-ray emission more than the SZ effect, because it is closer, but it is a relatively low density group with a larger scale N-S distribution not clearly related to the observed X-ray emission. We thus consider it unlikely that the foreground structure is responsible for the observed offset.  In a relaxed structure the shape of the gravitational potential close to the centre of the mass distribution will determine both the X-ray and the SZ effect distribution. Hence, we expect that the signal from the SZ effect and the free-free X-ray emission from the intracluster medium to peak in the same position. This means that we must have different distributions of electron density ($n_e$) and temperature ($T_e$) in this region.  Since the weak lensing peak seems to be better align with X-ray rather than with SZ the displacement is most likely to be caused by an increased gas clumping factor at the position of the X-ray peak or a temperature excess in the position of the SZ peak.  The suggestion of a merging front between the NE and SW clumps (Maughan et al. 2003) might indicate a shocked higher density region, but the suggested location of this merging front is too far away to cause the 45 arcsec displacement. Furthermore if these two sub-clumps are merging for the first time it is hard to see how it could already strongly effect the region near the NE clump and far in front of the proposed merging front.  The galaxy distribution (Demarco et al. 2005) in the region of the SZ peak seems to indicate the presence of few galaxies (see Fig. \\ref{fig:opt_sz}). No spectra are currently available for them, but photometric analysis will be used to estimate their redshift. Those galaxies could also be part of the cluster, representing yet another merger component along that SE-NW axis. \\begin{figure} \\epsscale{0.8} \\plotone{f4.eps} \\caption{SZ effect contours as in Fig. \\ref{fig:sz} overlapped to the K-band IRIS2 image.}\\label{fig:opt_sz} \\end{figure}  A more likely scenario is that this displacement is a post merger effect due to the passage of a smaller clump (e.g. the one to the SE) through the denser NE clump. Maughan et al. (2006) suggested that the NE-SW merger happened more recently than a merger with the SE group, relying on the fact that the projected separation of the structures along the NE-SW direction is much smaller than the separation between NE and SW subclumps, and because the X-ray morphology appears more disrupted along this NE-SW axis in the Northern subcluster. However, the passage of the SE subclump through the NE one would destroy the equilibrium in the NE clump and could generate regions of shocked gas causing excess X-ray emission on one side (including also an elongation of the NE X-ray peak towards North, perpendicularly to the crossing direction). It would also heat the gas in the potential of the NE clump further enhancing the SZ effect relative to the X-ray emission in this region.  It is also plausible that there is a mixture of cooler gas associated with the cluster and much hotter gas [including possibly very hot $\\gtrsim20$ keV gas that is hard to detect with XMM or Chandra as shown for the `bullet' cluster by Markevitch et al. (2002)] associated with the merger, projected along the line of sight. Depending on the densities of these components, the X-ray and SZ effect would respond very differently to the two components, with the hotter components going undetected in the current X-ray data. This, along with the relatively poor X-ray statistics in this region could explain the lack of evidence for enhanced pressure at the SZE peak in the XMM data.  Deeper high resolution SZ (including the SW peak) and X-ray observations might help finding a conclusive description for the dynamics of this cluster.  \\subsection{The impact on cosmology} A clear consequence of our finding is that any cosmological conclusion based on the assumption that the SZ effect and the X-ray emission result from the same relaxed gas distribution could be wrong for clusters at high redshift like Cl~J0152-1357. Low resolution SZ effect data do not allow  the proper density and temperature distribution to be disentangled, with the risk that if the peak of the signal is blindly associated with the X-ray peak an error in the estimation of cosmological parameters is introduced.  The cosmological parameters are usually estimated by comparing the angular diameter distance $D_A$ as estimated by the X-ray and SZ effect. In order to determine the distance it is necessary to properly model the structure of the intracluster medium, to connect the quantities determied among integrations along the cluster line of sight with the cluster appearence on the plane of the sky: typically isothermal and spherical symmetric density distributions are assumed to extract $D_A$ and cosmological parameters. The failures of the assumptions for several reasons, including mergers, asphericity, clumpiness and nonisothermality induces errors on the cosmological parameters which increase with redshift: numerical simulation (Molnar, Birkinshaw \\& Mushotzky 2002, Yoshikawa, Itoh, and Suto 1998, Inagaki, Suginohara, \\& Suto 1995, Roettiger, Stone, \\& Mushotzky 1997, Mohr et al. 1999) have demonstrated that at $z\\sim 1$ the cumulative effect of isothermality and asphericity, almost negligible at low redshift but can introduce up to a 20\\% error in the estimation of the angular diameter distance and that the assumption of isothermality for merging subclusters overestimates $D_A$ by a 15-20\\%. The error introduced by this offset is likely to be much larger.  The error is larger the farther from virialization, hence clusters at higher redshift which exhibit less virialized structures than local ones could be more seriously affected. This also has implications on the studies of cluster evolution, because estimates of barionic mass and intracluster medium properties might be wrong.  In conclusion we stress the importance that only virialized clusters can be of use for deriving cosmological parameters. Furthermore, wherever X-ray observations indicate that the structure of the cluster is not yet relaxed the high resolution SZ effect data is a powerful tool for the analysis of the cluster properties. In combination with X-ray observations of free-free emission it may provide a powerful tool to study the effects of shocks and other phenomena which will be important for our understanding of the evolution of galaxy clusters."
    },
    "1005/1005.4276_arXiv.txt": {
        "abstract": "The spectral triple approach to noncommutative geometry allows one to develop the entire standard model (and supersymmetric extensions) of particle physics from a purely geometry stand point and thus treats both gravity and particle physics on the same footing. The bosonic sector of the theory contains a modification to Einstein-Hilbert gravity, involving a nonconformal coupling of curvature to the Higgs field and conformal Weyl term (in addition to a nondynamical topological term). In this paper we derive the weak field limit of this gravitational theory and show that the production and dynamics of gravitational waves are significantly altered. In particular, we show that the graviton contains a massive mode that alters the energy lost to gravitational radiation, in systems with evolving quadrupole moment. We explicitly calculate the general solution and apply it to systems with periodically varying quadrupole moments, focusing in particular on the the well know energy loss formula for circular binaries. ",
        "introduction": "\\label{intro} Noncommutative Geometry (NCG) is a gravitational theory which, even in its simplest form, can explain the Standard Model of particle physics, and account for all current experimental data, in a rather simple and certainly elegant way.  The simple --- in the sense that it generalizes the continuum Riemaniann manifold by considering its product by a discrete two points space --- NCG proposal should be certainly replaced by a less trivial noncommutative space as one reaches Planckian energy scales. Nevertheless, this is the proposal we have at hand, and given its success in accounting for the Standard Model of particle physics, it offers a valid theoretical framework to address early universe issues. Alternatively, one can use experimental high energy physics data and astrophysical observations/measurements in order to test this NCG proposal and constrain its parameters. This is the approach used in this study.  One should indeed view this NCG proposal as an effective theory, which can however offer a valuable information about any NCG approach. In what follows we derive explicitly the weak field limit of this gravitational theory and then show that the production and dynamics of gravitational waves are both considerably modified from those obtained within the familiar General Relativity approach.  More precisely, in Section~\\ref{NCG-section} we give a short introduction to the noncommutative geometry spectral action, the framework within which we will then focus. In Section~\\ref{pert-equations} we first state in detail the conventions and signature we use and we analyze the issue of gauge conditions. We then analyze linear perturbations around a Minkowski background metric and we solve the noncommutative geometry gravitational wave equation in terms of the retarded Green's function.  We find that gravitational waves are only sourced from systems with a nontrivial quadrupole moment, as within General Relativity, while the NCG theory contains massive as well as massless gravitons. In Section~\\ref{examples} we concentrate first on some simple and then on some physical examples. Using the requirement that the mass of the gravitons must be positive (and real), we can fix the sign of the couplings in the NCG spectral action approach.  We then calculate the energy loss for a circular binary system and compare it to the results obtained from standard General Relativity. We conclude that the amplitude of modifications within NCG is small, nevertheless the NCG approach leads to some distinctive features which we analyze.  We round up with our conclusions in Section~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions} NonCommutative Geometry is a natural extension to our familiar notions of Riemannian geometry, that has the additional benefit of producing the action of all the Standard Model fields in addition to gravity terms, purely through geometrical considerations. Thus NCG treats both gravity and matter on an equal footing and provides us with concrete relationships between matter and gravitational couplings. The gravitational sector of (the asymptotic expansion of) this theory produces modifications to General Relativity and in this paper we explore the ramifications of these modifications on the formation and evolution of gravitational waves.  We have shown that the theory contains both massive and massless gravitons and that the requirement that the mass of these gravitons be positive fixes the sign of one of the couplings in the theory (for a given choice of sign conventions). We also show that both these modes are sourced by the quadrupole moment of a system (just as in standard GR) and that the retarded Green's function is not restricted to the past light cone of the observer (unlike GR), as one would expect for a system with massive modes. We have explicitly calculated the energy loss for a circular binary system and compared the results to those of standard GR. We have demonstrated that the amplitude of these NCG modifications is suppressed by the distance between the observer and the source of the gravitational waves and hence will typically be small.   Despite the extremely small amplitude of deviations from standard results, we have shown that NCG produces several distinctive features. Firstly, the amplitude of the energy lost by a binary pair can be higher or lower than the expected value, depending on the orbital period of the pair and the distance to the observer. This opens up the possibility that the observed energy loss from such a pair would be seen to oscillate as the binary moves with respect to the Earth. Whilst such effects are likely to be beyond current observational resolution, they allow for an unexpected beat phenomenon, which would be a concrete signature of NCG.  In addition, we have shown that the amplitude of these effects is (approximately) proportional to $(1-2\\omega/c\\beta)^{-1}$, where $\\omega$ is the orbital frequency of the binary. Thus, it would appear that the NCG corrections to the energy loss by the binary can become arbitrarily large as the frequency of the binary approach the critical frequency $\\beta c$. In such a regime, the weak field approach taken here would no longer be valid (and numerical approximations break down), so one would not trust systems very close to this limit, however it is certainly true that astrophysical constraints on the parameters of the theory will be significantly improved for objects with a very rapidly changing quadrupole moment. A precise understanding of such systems is likely to require detailed knowledge of various astrophysical effects (radiation and particle production, tidal stripping etc.) as well as analytic solutions to the graviton field in the large frequency regime.  Finally, the form of Eq.~(\\ref{eq:grav_pert}) suggests that similar behaviour may be present in other systems, with periodic, or almost periodic, variations in the (mass) quadrupole moment.  For laboratory systems, the gravitational radiation predicted by General Relativity is negligible, however if the NCG enhancement were sufficiently large, this may no longer be true. Laboratory systems regularly have very large oscillation frequencies ({\\rm e.g.}, lattice vibrations in solids can easily exceed $10^{12} {\\rm Hz}$) which would experience anomalous damping, if the system was producing significant amounts of gravitational radiation. This opens up the (remote) possibility that the noncommutative nature of space-time might be probed in the laboratory.  One can immediately use the results of this paper to examine circular binary systems, in order to constrain the value of $\\beta$~\\cite{wjm}. Similarly, one can include eccentricity which may result in more restrictive constraints on the theory. An alternative avenue would be to use the gravitational wave-forms given here to deduce the consequences for direct gravity wave searches (LIGO, VIRGO, LISA, {\\sl etc}).  In particular, to extend these result to the large field regime and look for modifications to the {\\sl chirp} that develops at the end of in-spiral events.   \\vskip1truecm {\\bf Acknoweldgments} \\medskip \\\\ The work of W.\\ N.\\ is supported in part by the NSF grant PHY0854743, the George A. and Margaret M. Downsbrough Endowment and the Eberly research funds of Penn State.  The work of M.\\ S.\\ is partially supported by the European Union through the Marie Curie Research and Training Network {\\sl UniverseNet} (MRTN-CT-2006-035863). J.\\ O.\\ acknowledges support from the Alfred P. Sloan Foundation and the Eberly College of Science."
    },
    "1005/1005.0275_arXiv.txt": {
        "abstract": "% {CoRoT and Kepler missions are now providing high-quality asteroseismic data for a large number of stars. Among intermediate-mass and massive stars, fast rotators are common objects. Taking the rotation effects into account is needed to correctly understand, identify, and interpret the observed oscillation frequencies of these stars. A classical approach is to consider the rotation as a perturbation.}% {In this paper, we focus on gravity modes, such as those occurring in $\\gamma$~Doradus, slowly pulsating B (SPB), or Be stars. We aim to define the suitability of perturbative methods.}% {With the two-dimensional oscillation program (TOP), we performed complete computations of gravity modes -- including the Coriolis force, the centrifugal distortion, and compressible effects -- in 2-D distorted polytropic models of stars. We started with the modes $\\ell=1$, $n=1$--14, and $\\ell=2$--3, $n=1$--5, 16--20 of a nonrotating star, and followed these modes by increasing the rotation rate up to 70\\% of the break-up rotation rate. We then derived perturbative coefficients and determined the domains of validity of the perturbative methods.}% {Second-order perturbative methods are suited to computing low-order, low-degree mode frequencies up to rotation speeds $\\sim$100\\unit{km\\,s^{-1}} for typical $\\gamma$~Dor stars or  $\\sim$150\\unit{km\\,s^{-1}} for B stars. The domains of validity can be extended by a few tens of \\unit{km\\,s^{-1}} thanks to the third-order terms. For higher order modes, the domains of validity are noticeably reduced. Moreover, perturbative methods are inefficient for modes with frequencies lower than the Coriolis frequency $2\\Omega$. We interpret this failure as a consequence of  a modification in the shape of the resonant cavity that is not taken into account in the perturbative approach.}% {} ",
        "introduction": "\\label{Sec:Intro}  CoRoT \\citep[Convection, Rotation and planetary Transits,][]{BaglinM06_COROT} and Kepler \\citep{BoruckiK07_KEPLER} are space missions providing uninterrupted high-quality photometry time series over several months or years ideally suited for asteroseismic study. Asteroseismology provides very accurate determinations of the stellar parameters (mass, radius, age, etc.) and probes stellar structure to constrain physical processes occurring in stars. The first step towards this goal requires correctly understanding the structure of the observed oscillation spectra, and especially correctly identifying the observed modes. In the case of main-sequence \\citep[e.g.][]{Michel08,Benomar09,Chaplin10} and giant  \\citep[e.g.][]{Miglio09,Hekker09,Bedding10} FGK stars with solar-like oscillations, the spectrum structure is well understood, which eases interpretation.  The spectra of classical pulsators is often noticeably more complex. For instance, the high-quality observations of $\\delta$ Scuti \\citep[e.g.][]{GarciaHernandez09,Poretti09} and $\\gamma$~Doradus \\citep[e.g.][]{Mathias09} stars have exhibited very rich and complex spectra of acoustic (p) and gravity (g) modes, respectively, containing several hundred --or more-- modes. Interpreting their spectra is very challenging today. Indeed, these stars generally spin rapidly, so the effects of rotation on the mode frequencies must be considered.  Here, we are concerned with gravity modes, i.e., low-frequency modes driven by the buoyancy force. They are excited and observed in a broad panel of stars, for instance, in $\\gamma$~Dor, SPB and some Be stars. The $\\gamma$~Dor stars form a class of main-sequence stars with type around F0V that can sometimes rotate rapidly \\citep[e.g.][and reference therein]{DeCat06}, while the rotation rate of Be stars is extreme, usually very close to their break-up limit $\\Omega\\approx\\Omega_K\\equiv\\sqrt{GM/R^3}$ \\citep[e.g.][]{Fremat06}.  The effects of rotation on the oscillation modes can be treated as a perturbation where the rotation rate is the small parameter. A 1st-order correction has been proposed by \\citet{Ledoux51}, 2nd-order by \\citet{Saio81}, \\citet{DziembowskiGoode92}, or \\citet{Suarez06},  and 3rd-order terms have been developed by \\citet{Soufi98}. While perturbative methods are expected to be accurate enough for slowly rotating stars, their true domain of validity cannot be determined in the absence of exact calculations to compare them with.  In the past few years, calculations of p~modes with both the centrifugal distortion and the Coriolis force have been performed in polytropic models of stars \\citep{Lignieres06,Reese06} and realistic 2-D stellar structures \\citep{Lovekin08, Reese09}. \\citet{Lignieres06} and \\citet{Reese06} have shown that, above $\\Omega \\sim 0.15\\Omega_K$, perturbation methods fail to reproduce low-degree and low-order p-mode frequencies ($\\ell \\le 3$ and $n \\le 10$) with the accuracy of CoRoT long runs. The structure of the modes is also drastically modified, and this leads to deep changes in the structure of the p-mode spectrum \\citep{Reese08,Lignieres08,Lignieres09}.  We used an oscillation code based on \\citet{Reese06} to perform g-mode calculations with a complete description of the rotational effects on the modes. In this paper, we focus on the limits of validity for perturbative methods. The models and the method are described in Sect.~\\ref{Sec:Methods}. We then derive the perturbative coefficients from the complete computations (Sect.~\\ref{Sec:Perturbative}), and compare the results obtained with both methods to determine and discuss the domains of validity for perturbative methods (Sect.~\\ref{Sec:Validity}) before concluding in Sect.~\\ref{Sec:Concl}. ",
        "conclusions": "\\label{Sec:Concl}  In the present work, we have computed accurate frequencies for g modes in polytropic models of uniformly spinning stars. We started from high- and low-frequency, low-degree ($\\ell \\le 3$) g modes of a nonrotating star and followed them up to $\\Omega = 0.7 \\Omega_K$. This allowed us to provide a table of numerically-computed perturbative coefficients up to the 3rd order for a polytropic stellar structure (with index $\\mu=3$). This table can serve as a reference for testing the implementation of perturbative methods. We then determined the domains of validity of perturbative approximations. For the high-frequency (low-order) modes, 2nd-order perturbative methods correctly describe modes up to $\\sim 100\\unit{km\\,s^{-1}}$ for $\\gamma$~Dor stars and up to $\\sim 150\\unit{km\\,s^{-1}}$ for B stars. The domains of validity can be extended by a few tens of \\unit{km\\,s^{-1}} with 3rd-order terms. However, the domains of validity shrink at low frequency. In particular, perturbative methods fail in the inertial domain $\\omega < 2\\Omega$ because of a modification in the shape of the resonant cavity.   In a next step, we plan to compare our complete computations with the so-called traditional approximation, which is also extensively used to determine g-mode frequencies \\citep[e.g][]{Berthomieu78,LeeSaio97}. We will also analyze how rotation affects the regularities of the spectrum  -- such as the period spacing -- and compare it to the predictions of the perturbative and traditional methods. In the present study, we have focused on low-degree modes, but a more complete exploration clearly needs to be performed. In particular, we might look for the singular modes predicted by \\citet{Dintrans00}. It requires to take care of dissipative processes, which play an important role in this case."
    },
    "1005/1005.2877_arXiv.txt": {
        "abstract": "{ We present a full high resolution SPIRE FTS spectrum of the nearby ultraluminous infrared galaxy $\\Mrk{231}$. In total 25 lines are detected, including CO $J=5{-}4$ through $J=13{-}12$, 7 rotational lines of $\\HtO$, 3 of $\\OHp$ and one line each of $\\HtOp$, $\\CHp$, and HF\\null.  We find that the excitation of the CO rotational levels up to $J=8$ can be accounted for by UV radiation from star formation.  However, the approximately flat luminosity distribution of the CO lines over the rotational ladder above $J=8$ requires the presence of a separate source of excitation for the highest CO lines. We explore X-ray heating by the accreting supermassive black hole in $\\Mrk{231}$ as a source of excitation for these lines, and find that it can reproduce the observed luminosities.  We also consider a model with dense gas in a strong UV radiation field to produce the highest CO lines, but find that this model strongly overpredicts the hot dust mass in $\\Mrk{231}$.  Our favoured model consists of a star forming disk of radius $560\\pc$, containing clumps of dense gas exposed to strong UV radiation, dominating the emission of CO lines up to $J=8$.  X-rays from the accreting supermassive black hole in $\\Mrk{231}$ dominate the excitation and chemistry of the inner disk out to a radius of $160\\pc$, consistent with the X-ray power of the AGN in $\\Mrk{231}$.  The extraordinary luminosity of the $\\OHp$ and $\\HtOp$ lines reveals the signature of X-ray driven excitation and chemistry in this region.} ",
        "introduction": "Carbon monoxide (CO) is a fundamental tracer of interstellar molecular gas. However, since only the lowest 3 rotational transitions are relatively easily accessible with ground-based telescopes, the diagnostic use of higher rotational levels is poorly developed. This hiatus in our knowledge is becoming acute now that high-$J$ CO observations of high-$z$ galaxies are becoming possible.  \\begin{figure*} \\centering \\includegraphics[width=\\textwidth]{14682fg1.eps} \\caption{SPIRE FTS spectrum of $\\Mrk{231}$. Line identifications are given in red for CO lines, in blue for $\\HtO$, in magenta for $\\OHp$, in cyan for $\\HtOp$, and in green for the remaining lines. } \\label{fig.spec} \\end{figure*}  With this in mind, we have embarked on the {\\it Herschel\\/} Comprehensive (U)LIRG Emission Survey (HerCULES), an open time Key Project on the ESA {\\it Herschel\\/} Space Observatory\\footnote{ {\\it Herschel\\/} is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA.}  (Pilbratt et al., this issue).  The HerCULES project will establish a comprehensive inventory of the gas cooling lines in a flux-limited sample of 29 (Ultra)luminous Infrared Galaxies or (U)LIRGs, using high spectral resolution observations with the Fourier Transform Spectrograph (FTS) of the SPIRE instrument (Griffin et al., this issue), combined with observations of the [$\\CII$] $158\\mum$ line and the [$\\OI$] 63 and $146\\mum$ lines with PACS (Poglitsch et al., this issue). Key aims of HerCULES are the development of the diagnostic use of the gas cooling lines in local (U)LIRGs, and establishing a local benchmark for observations of high-$z$ galaxies with the Atacama Large Millimeter Array.  In addition, since the FTS yields full spectra, any other luminous emission lines detected (e.g., of $\\HtO$) will be available for study.  \\citet{SpaansMeijerink08} have shown that X-ray excitation of the gas (e.g., by an AGN) and UV irradiation by young massive stars produce very different luminosity distributions over the CO rotational lines. Physically, the difference arises because X-rays penetrate a larger column density of gas than UV photons, and are less effective in dissociating the molecules. In addition, while the gas heating efficiency in a photon dominated region (PDR) is less than $1\\%$, in X-ray dominated regions (XDRs) this efficiency is $10{-}50\\%$.  As a result, for comparable irradiated energies, XDRs tend to have larger column densities of warmer molecular gas than PDRs, and will produce much more luminous emission in the high-$J$ CO lines. In contrast, PDRs are more efficient than XDRs in heating the dust. X-rays also give rise to significant ionization in the molecular gas and therefore drive an efficient ion-molecule chemistry, leading to pronounced chemical differences between PDRs and XDRs \\citep{MeijerinkSpaans05}. Testing and using these diagnostics is one of the principal aims of the HerCULES project.  In this Letter, we discuss the first results of the HerCULES programme and present the SPIRE FTS spectrum of the nearby ULIRG $\\Mrk{231}$ ($\\sou{UGC}{8058}$, $\\IRAS{F12540{+}5708}$), the most luminous galaxy in the Revised IRAS Bright Galaxy Sample \\citep{Sandersetal03}.  Adopting $z = 0.042170$ as the heliocentric redshift of $\\Mrk{231}$, correcting for the local flow, and applying a flat 5-year WMAP cosmology ($H_0 = 70\\kms\\per{Mpc}$, $\\Omega_{\\Lambda} = 0.73$) yields a luminosity distance $\\qu{D}{L}=192\\Mpc$, with $1''$ corresponding to $0.856\\kpc$, as provided by NED\\footnote{\\tt http://nedwww.ipac.caltech.edu/}.  The derived $8-1000\\mum$ luminosity of $\\Mrk{231}$ is then $\\LIR=4.0\\cdot10^{12}\\Lsun$. $\\Mrk{231}$ contains a luminous, optically visible AGN, classified as a Seyfert~1 or a Broad Absorption Line QSO \\citep{Boksenbergetal77}.  A highly absorbed power-law X-ray spectrum was observed by \\citet{Braitoetal04} with $\\qu{L}{X}=6^{+0.6}_{-0.3}\\cdot 10^{43}\\un{erg}\\ps$ between 2 and $10\\un{keV}$. However, $\\Mrk{231}$ also contains a kpc size disk harbouring intense star formation as shown by high resolution radio imaging \\citep{Tayloretal99}. Interferometric imaging of CO $J=1{-}0$ and $J=2{-}1$ emission shows an inner disk of radius $\\sim520\\pc$, containing $45\\%$ of the total molecular gas mass, and embedded in a more extended and diffuse emission component; the total molecular gas mass is $5\\cdot10^9\\Msun$ \\citep{DownesSolomon98}. The lowest part of the CO ladder (up to $J=6{-}5$), was analysed by \\citet{Papadopoulosetal07}, who showed that the integrated CO emission can provide a significant contribution to the total gas cooling. Indications for X-ray-driven chemistry have been found by \\citet{Aaltoetal07} in HNC and HCN line ratios, and by \\citet{GonzalezAlfonsoetal08} in the abundances of OH and $\\HtO$ observed in absorption with ISO\\null. In a comprehensive study of ULIRGs and low-$z$ quasars, \\citet{Veilleuxetal09} derive a fraction of $70{\\pm}15\\%$ for the AGN contribution to the far-infrared luminosity of $\\Mrk{231}$, with the remainder coming from star formation. ",
        "conclusions": "\\subsection{CO excitation} \\label{sec.PDRXDR}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{14682fg2.eps}} \\caption{Luminosities of CO lines from $\\Mrk{231}$. Filled symbols represent measurements from the SPIRE FTS spectrum, while ground-based measurements are denoted with open symbols. Coloured lines indicate two model PDR components (red and green lines) and an XDR component (blue line). The sum of these three components is indicated by the black line and fits the CO measurements. In the legend, $n$ denotes the number density of hydrogen nuclei ($n=\\qu{n}{H}+2\\nHt$) in cm$^{-3}$, $\\qu{G}{0}$ denotes the incident UV flux in units of $1.6\\cdot10^{-3}\\esc$ for the PDRs, and $\\qu{F}{X}$ the incident X-ray flux for the XDR\\null. The legend also indicates the relative emitting areas of the three components.} \\label{fig.CO} \\end{figure}  We combine the CO line fluxes from the spectrum shown in \\figref{fig.spec} with ground-based measurements of the lower lines \\citep[][and references therein]{Papadopoulosetal07} in order to construct the CO rotational excitation diagram shown in \\figref{fig.CO}. It is seen that an approximately flat luminosity distribution is obtained for the lines from $J=5$ upwards. Note that the CO $J=10{-}9$ line is blended with the $\\HtO$ $3_{1,2}{-}2_{2,1}$ line which is expected to have some luminosity ($1.8\\cdot 10^7\\Lsun$ for the model by Gonz\\'alez-Alfonso et al., this issue), and this may account for its somewhat high flux. The total luminosity measured in the CO lines up to $J=13{-}12$ is $(2.8\\pm1.1)\\cdot10^8\\Lsun$. Note that only $4\\%$ of the total CO line luminosity is contained in the lowest 3 transitions. For comparison, in our Milky Way this fraction is $43\\%$ \\citep{Fixsenetal99}.  The approximately flat distribution of CO line luminosity with rotational level indicates that several excitation components must be present, since individual components always produce a more peaked excitation diagram. We model these components using the one dimensional PDR/XDR models of \\citet{Meijerinketal07}, as shown in \\figref{fig.CO}. The CO lines up to $J=8{-}7$ can be produced by a combination of 2 PDRs, in qualitative agreement with the decomposition by \\citet{Papadopoulosetal07}. However, a challenge is presented by the highest CO lines: $J=13{-}12$ and $J=12{-}11$, arising from energy levels 503 and $461\\K$ above the ground state. As shown in \\figref{fig.CO}, these lines are strongly underproduced by the PDRs dominating the emission in the lower lines, since the resulting gas temperatures are not high enough for significant population of the $J>10$ levels. These lines therefore require the presence of a third excitation component, which can be either an XDR or a high excitation PDR\\null.  A model fit with an XDR producing the highest CO lines is shown in \\figref{fig.CO}. The required X-ray illumination for this XDR can be produced by the AGN in $\\Mrk{231}$ \\citep{Braitoetal04}, out to a distance of $160\\pc$ from the nucleus, ignoring absorption. The ratio of radiating surfaces in the model shown in \\figref{fig.CO} implies an extended low excitation PDR component (green curve), with a less extended and denser central XDR region (blue curve). Dense clouds with a smaller surface, close to massive stars and probably embedded in the more diffuse component, account for the medium excitation component (red curve).  Alternatively, a very dense, high illumination PDR can account for the highest CO lines. A good fit is found with $n=10^{6.5}\\pcmcub$ and $\\qu{G}{0}=10^5$ and a surface ratio from medium to high excitation of $1.0:0.03$.  Here the small surface area for the high excitation PDR indicates a number of small high density clumps in a very strong UV field. Since the radiating surface of the high excitation PDR is about $30\\times$ smaller than that of the medium excitation PDR, but its density about $30\\times$ larger, the $\\Ht$ masses in these two components must be comparable.  For an O5 star, the required $G_0=10^5$ is reached at a radius of $0.3\\pc$. For a star formation rate of $100\\Msun\\pyr$ and a power law initial mass function with slope $-2.35$ between masses of 0.3 and $120\\Msun$, there are $7.6\\cdot10^5$ stars of spectral type O5 or earlier in $\\Mrk{231}$.  The total volume with $G_0>10^5$ is then $8.6\\cdot10^4\\pun{pc}{3}$, while (for a $520\\pc$ radius disk with $15\\pc$ thickness, following \\citet{Daviesetal04}) the total volume of the gas disk is {\\bf $1.3\\cdot10^7\\pun{pc}{3}$}.  In other words, in this scenario approximately half of the molecular mass would have to be contained in 0.7\\% of the total volume, and located within $0.3\\pc$ from an O5 (or hotter) star. Efficient UV heating by the $G_0=10^5$ radiation field would heat the dust in these clumps (in total $\\sim50\\%$ of the total dust mass in $\\Mrk{231}$) to a temperature of about $170\\K$. In contrast, in the XDR model, where dust heating would be less efficient, the dust temperature would only be $\\sim70\\K$ \\citep{MeijerinkSpaans05}.  These predictions can be tested by analysing the infrared SED of $\\Mrk{231}$. Gonz\\'alez-Alfonso et al.\\ (this issue), found that the hot ($T=150{-}400\\K$) component in the SED of $\\Mrk{231}$ accounts for about $\\sim20\\%$ of the total infrared luminosity, but is produced by only 0.02\\% of the total dust mass. This result limits the fraction of gas within $0.3\\pc$ from an O5 star in $\\Mrk{231}$ to much less than the $\\sim50\\%$ required to produce the highest CO lines with a high excitation PDR\\null. This problem does not exist for the XDR model, which predicts most of the dust to be cooler.  \\subsection{XDR chemistry}  The extraordinarily luminous emission from the molecular ions $\\HtOp$ and $\\OHp$ reveals the chemical signature of an XDR\\null.  Assuming that the emission arises from a disk with $160\\pc$ radius (as derived above), we can derive column densities in the upper levels of the relevant transitions, which results in values $\\qu{N}{up}\\sim3.0{-}3.5\\cdot10^{13}\\pun{cm}{-2}$ for both species.  Modeling the nuclear molecular gas disk in $\\Mrk{231}$ with a radius of $520\\kpc$ and $\\Ht$ mass of $2.2\\cdot10^9\\Msun$ then results in lower limits to the total $\\OHp$ and $\\HtOp$ abundances relative to $\\Ht$ of $\\sim2\\cdot10^{-10}$ in the central $160\\pc$. Given the short radiative lifetimes ($<60\\sec$) of the upper levels involved, most $\\HtOp$ and $\\OHp$ molecules will be in the ground state, and total abundances will exceed these lower limits by large factors.  Such abundances require an efficient and penetrative source of ionization in the molecular gas, since the production of $\\OHp$ is mainly driven by $\\rm H^+ + O \\rightarrow O^+ + H$ followed by $\\rm O^+ + H_2 \\rightarrow OH^+ + H$ and $\\rm H^+ + OH \\rightarrow OH^+ + H$. In the models considered in \\secref{sec.PDRXDR}, the key species $\\Hp$ is $2{-}3$ orders of magnitude more abundant in the XDR model than in the high excitation PDR, and the $\\OHp$ abundance is larger by a comparable amount.  A similar argument can be made for $\\HtOp$, which is formed by $\\rm OH^+ + H_2 \\rightarrow H_2O^+ + H$.  The extraordinary luminosity of the $\\OHp$ and $\\HtOp$ (and to a lesser extent $\\CHp$) lines in $\\Mrk{231}$ is underlined by a comparison with the SPIRE spectrum of the Orion bar PDR (Habart et al., this issue), which shows no trace of $\\OHp$ or $\\HtOp$, and only weak $\\CHp$ emission, while in $\\Mrk{231}$ the lines are only a factor $2{-}3$ fainter than the CO lines. While enhanced cosmic ray fluxes in a starburst environment will increase the degree of ionization and hence the production of $\\OHp$ and $\\HtOp$ in a PDR, they do not elevate the temperatures to the level required to produce the highest CO lines \\citep{Meijerinketal06}. It is thus the combination of strong high-$J$ CO lines and high $\\OHp$ and $\\HtOp$ abundances that reveals X-ray driven excitation and chemistry in $\\Mrk{231}$."
    },
    "1005/1005.5499_arXiv.txt": {
        "abstract": "{The redshift range from 2.2 to 3, is known as the 'redshift desert' of quasars because quasars with redshift in this range have similar optical colors as normal stars and are thus difficult to be found in optical sky surveys. A quasar candidate, SDSS J085543.40-001517.7, which was selected by a recently proposed criterion involving near-IR $Y-K$ and optical $g-z$ colors, was identified spectroscopically as a new quasar with redshift of 2.427 by the LAMOST commissioning observation in December 2009 and confirmed by the observation made with the NAOC/Xinglong 2.16m telescope in March 2010. This quasar was not targeted in the SDSS spectroscopic survey because it locates in the stellar locus of the optical color-color diagrams, while it is clearly separated from stars in the $Y-K$ vs. $g-z$ diagram. Comparing with other SDSS quasars we found this new quasar with $i$ magnitude of 16.44 is apparently the brightest one in the redshift range from 2.3 to 2.7. From the spectral properties we derived its central black hole mass as  $(1.4\\sim3.9) \\times 10^{10} M_\\odot$ and the bolometric luminosity as $3.7\\times 10^{48}$ \\ergs, which indicates that this new quasar is intrinsically very bright and belongs to the most luminous quasars in the universe. Our identification supports that quasars in the redshift desert can be found by the quasar selection criterion involving the near-IR colors. More missing quasars are expected to be recovered by the future LAMOST spectroscopic surveys, which is important to the study of the cosmological evolution of quasars at redshift higher than 2.2. ",
        "introduction": "% \\label{sect:intro}  After the discovery of first quasar (Schmidt 1963), many quasar surveys have been carried out in optical band and the number of quasars has increased steadily in the past four decades (Richards et al. 2009). Especially, a large number of quasars have been discovered in two recent spectroscopic surveys, namely, the Two-Degree Fields (2DF) survey (Boyle et al. 2000) and Sloan Digital Sky Survey (SDSS) (York et al. 2000). 2DF has discovered more than 20,000 quasars (Croom et al. 2004), and SDSS has identified more than 100,000 quasars (Schneider et al. 2010; Abazajian et al. 2009). 2DF mainly selected lower redshift ($z<2.2$) quasars with UV-excess (Smith et al. 2005), while SDSS adopted a multi-band optical color selection method for quasars mainly by excluding the point sources in the stellar locus of color-color diagrams (Richards et al. 2002).  Some dedicated methods were also proposed for finding high redshift quasars (Fan et al. 2001a,b; Richards et al. 2002). However, the efficiency of identifying quasars with redshift between 2.2 and 3 is obviously low in SDSS (Schneider et al. 2010), because quasars with such redshift usually have similar optical colors as stars and are thus mostly excluded by the SDSS quasar candidate selection algorithm. Therefore, the redshift range from 2.2 to 3 is often regarded as the `redshift desert' of quasars because of the difficulty in identifying quasars within this redshift range.  In addition, the low efficiency of finding quasars with redshift between 2.2 and 3 has led to the obvious incompleteness of quasar sample in this redshift range and serious problems in constructing the luminosity function for quasars. More importantly, many studies have shown that the quasar activity actually peaks at redshift range $2<z<3$ (see Richards et al. 2006; Jiang et al. 2006). Therefore, recovering the missing quasars with redshift between 2.2 and 3 has become an important task in the quasar study.  Although quasars in the redshift desert have similar optical colors as stars, they are usually more luminous than normal stars in the infrared K-band because the fluxes of normal stars decreases rapidly in the near-IR bands while quasar SEDs are relatively flat (Warren et al. 2000). An important way of finding these missing quasars has been suggested by using the infrared K-band excess based on the UKIRT (UK Infrared Telescope) Infrared Deep Sky Survey (UKIDSS) (Warren et al. 2000; Hewett et al. 2006; Maddox et al. 2008). Combining the UKIDSS $YJHK$ and SDSS $ugriz$ magnitudes, some criteria to separate quasars and stars have been proposed previously. Maddox et al. (2008) suggested a selection criterion of $z<4$ quasar candidates in the $g-J$ vs. $J-K$ diagram. Chiu et al. (2007) investigated the different color-color diagrams in optical and near-IR bands with a sample of 2837 SDSS-UKIDSS quasars, and found that the $g-r$ vs. $u-g$ diagram and the $H-K$ vs. $J-H$ diagram are more effective in separating quasars and stars than other diagrams. They also proposed to use the $Y-K$ vs. $u-z$ diagram to select low redshift ($z<3$) quasars. Recently, based on a SDSS-UKIDSS sample of 8498 quasars, Wu \\& Jia (2010) proposed to use the $Y-K$ vs. $g-z$ diagram to select $z<4$ quasars and use the $J-K$ vs.$i-Y$ diagram to select $z<5$ quasars. Although with these two criteria we can recover 8447 of 8498 SDSS-UKIDSS quasars (with a percentage of 99.4\\%), we still need to demonstrate whether we can efficiently discover new quasars, especially those in the redshift desert, by applying our criteria to select quasar candidates in the SDSS spectroscopically surveyed area.  The Large Sky Area Multi-Object Fibre Spectroscopic Telescope (LAMOST) is a 4-meter reflecting Schmidt telescope with 20 square degree field of view (FOV) and 4000 fibers in the focal plane (Su et al. 1998), located in the NAOC/Xinglong station. After finishing its main construction in 2008, LAMOST has entered the commissioning phase since 2009. Some test observations have been done in the winter of 2009. Although LAMOST has not reach its full capability in the commissioning phase, these observations already led to the discovery of some new quasars, including a bright quasar with redshift of 2.427, which is the first quasar in the `redshift desert' discovered by LAMOST.   \\begin{figure} \\centering \\includegraphics[width=11.0cm, angle=0]{fig1.eps} \\caption{The finding chart of SDSS J085543.40-001517.7. The size is 200\"$\\times$200\". } \\label{Fig1} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=14.0cm, angle=0]{fig2.eps} \\caption{The location of SDSS J085543.40-001517.7 (blue triangle) in three optical color-color diagrams (a,b,c) and the $Y-K$ vs. $g-z$ diagram (d), in comparing with the 8996 SDSS-UKIDSS stars. Black and red dots represent the normal and later type stars, respectively. Dashed line shows the $z<4$ quasar selection criterion proposed by Wu \\& Jia (2010).} \\label{Fig2} \\end{figure} ",
        "conclusions": "\\label{sect:discussion}  Quasars with redshifts in the range from 2.2 to 3  are very important for studying their cosmological evolution, and the relation between quasar activity and star formation activity which peaks at redshift between 1 and 2 (Madau et al. 1998). However, because these quasars have similar optical colors as normal stars, it is very difficult for find them in previous quasar surveys. The low efficiency of finding quasars in the `redshift desert' has led to the obvious incompleteness of quasar sample in this redshift range and serious problems in constructing the luminosity function for quasars around the redshift peak (between 2 and 3) of quasar activity (Richards et al. 2006; Jiang et al. 2006).  In this paper we have presented a case study to find a very bright new quasar in the redshift desert by the LAMOST commissioning observation. The spectroscopic identification of an $i=16.44$ source, SDSS J085543.40-001517.7, as a $z=2.427$ quasar gives us confidence to discover more missing quasars in the future LAMOST quasar survey. This discovery also supports the idea that by combining the UKIDSS near-IR colors with the SDSS optical colors we are able to efficiently recover these missing quasars. In the winter of 2009, LAMOST has made test observations on several sky fields and we are now searching for more quasars from the spectra taken in these fields. The discovery of more new quasars in these fields will be reported in the future works. We hope that in the next a few months great progress will be made in improving the capability of LAMOST spectroscopy and the spectral processing pipeline. As long as LAMOST can reach its designed capability after the commissioning phase, we expect to find several hundred-thousands of quasars in the LAMOST quasar survey. This will form the largest quasar sample in the world and play a leading role in the quasar study of the next decade.     \\normalem"
    },
    "1005/1005.2046_arXiv.txt": {
        "abstract": "We compute here the spin-orbit and spin-spin couplings needed for an accurate computation of the phasing of gravitational waves emitted by comparable-mass binaries on eccentric orbits at the second post-Newtonian (PN) order. We use a quasi-Keplerian parametrization of the orbit free of divergencies in the zero eccentricity limit. We find that spin-spin couplings induce a residual eccentricity for coalescing binaries at 2PN, of the order of $10^{-4}$-$10^{-3}$ for supermassive black hole binaries in the LISA band. Spin-orbit precession also induces a non-trivial pattern in the evolution of the eccentricity, which could help to reduce the errors on the determination of the eccentricity and spins in a gravitational wave measurement. ",
        "introduction": "Up to now most studies concerning gravitational wave emission from binaries have been done assuming circular orbits (see e.g.~\\cite{petf,kjs} and references therein). Numerical studies of the formation of binaries and of their subsequent development suggest instead that the orbits could be eccentric~\\cite{caab,bpbms}. It is thus of relevance to take the eccentricity into account when one computes the waveform emitted by such a system. This problem has not yet been fully explored, mainly due to its complexity, and here as a further step we derive the equations which govern the time evolution of the orbital parameters, and in particular the eccentricity, including the spin-orbit and the spin-spin couplings needed for an accurate computation in the post-Newtonian approximation up to 2PN of the phasing of the gravitational waves emitted by binary systems, with components of comparable mass.  The PN approximation is valid when the two object forming the binary are sufficiently separated. The issue of estimating the limit of its validity has been tackled with different methods: comparisons between post-Newtonian templates and results from numerical relativity, for non-spinning binaries on quasi-circular orbits~\\cite{bcgshhb}; and comparisons between different post-Newtonian waveforms for spinning and non-spinning binaries on quasi-circular orbits~\\cite{bcv1,bcv2}. All of these studies found a remarkable reliability of the PN approximation up to separations as small as the innermost stable circular orbit, $r=6GM/c^2$. It is not clear if that still holds for eccentric binaries, and this must be answered by extending these comparisons to such systems, but one can trust that for low enough eccentricities, the post-Newtonian approximation is reliable up to the end of the inspiral.  We also derive a quasi-Keplerian parametrization of the orbit free of divergencies in the zero eccentricity limit, and find that spin-spin couplings induce a residual eccentricity of 2PN order after the orbit has been circularized by gravitational wave emission.  We then solve the equations which govern the evolution of the eccentricity and the mean motion for different values of the masses and spins as a function of the initial eccentricity. ",
        "conclusions": "The main result of this paper is the derivation of the spin-spin effects in the evolution of the mean motion and of the eccentricity, for binaries with an arbitrary eccenticity. Particularly, the fact that spin-spin couplings may induce a residual eccentricity can be important for parameter estimation when gravitational wave observations will be made possible. If eccentric templates allow to measure eccentricities of $O(10^{-3}$-$10^{-4})$, the modulation induced by spin-orbit precession could significantly improve the determination of the spins of the binary.  We also derived the equations of motion $r(t)$ and $\\phi(t)$, for black hole binaries of comparable mass with Newtonian, spin-orbit, and spin-spin terms on eccentric orbits, and found a family of parametrizations free of divergencies in the circular limit $e \\to 0$."
    },
    "1005/1005.5720_arXiv.txt": {
        "abstract": "In recent years, analyses of eclipsing binary systems have unveiled differences between the observed fundamental properties of low-mass stars and those predicted by stellar structure models. Particularly, radius and effective temperatures computed from models are $\\sim 5-10$\\% lower and $\\sim 3-5$\\% higher than observed, respectively. These discrepancies have been attributed to different factors, notably to the high levels of magnetic activity present on these stars. In this paper, we test the effect of magnetic activity both on models and on the observational analysis of eclipsing binaries using a sample of such systems with accurate fundamental properties. Regarding stellar models, we have found that unrealistically high spot coverages need to be assumed to reproduce the observations. Tests considering metallicity effects and missing opacities on models indicate that these are not able to explain the radius discrepancies observed. With respect to the observations, we have tested the effect of several spot distributions on the light curve analysis. Our results show that spots cause systematic deviations on the stellar radii derived from light curve analysis when distributed mainly over the stellar poles. Assuming the existence of polar spots, overall agreement between models and observations is reached when $\\sim 35$\\% spot coverage is considered on stellar models. Such spot coverage induces a systematic deviation in the radius determination from the light curve analysis of $\\sim 3$\\% and is also compatible with the modulations observed on the light curves of these systems. Finally, we have found that the effect of activity or rotation on convective transport in partially radiative stars may also contribute to explain the differences seen in some of the systems with shorter orbital periods. ",
        "introduction": "\\label{sec:introduction}  Analyses of double-lined eclipsing binary stars (hereafter EBs) have provided fundamental properties of stars on the low-mass end of the Main Sequence. Such EB systems have revealed to be the best candidates to test stellar structure models since they can potentially yield masses and radii for M-type stars with uncertainties below $\\sim 1$\\% \\citep{Morales2009a}. However, results coming from a variety of independent studies indicate that differences exist between the observed properties and those predicted by models; for the same mass ($M$), models predict $\\sim 5-10$\\% smaller radius ($R$) and $\\sim 3-5$\\% larger effective temperatures (\\Teff) than observed for stars below 1~M$_{\\odot}$.  In contrast, the stellar luminosity for these binary components is in agreement with those of single stars and also with theoretical models \\citep[see][for a review]{Ribas2006}.  Several authors have argued that activity could be the cause of discrepancies found between models and observations of these systems \\citep{Torres2002,LopezMorales2005,Torres2006,LopezMorales2007,Ribas2008}. Further support to this hypothesis is lent by similar radius discrepancies found when comparing single magnetically active late-type stars with their inactive counterparts \\citep{Morales2008}, and by the results of interferometric observations of single inactive stars \\citep{Demory2009}. Until now, all the known low-mass EBs are close binaries with periods well below 10 days, thus they are expected to be synchronized \\citep{Mazeh2008}, i.e., with their rotation tidally locked to the orbital motion, and therefore they are fast rotators ($v_{\\rm rot} \\gtrsim 10$~km~s$^{-1}$). Such fast rotation in presence of magnetic fields could trigger high levels of activity that may modify the structure of these stars in some respects. Observationally, the high activity of these stars is usually confirmed by saturated X-ray emission ($\\log L_{X}/L_{bol}\\sim-$3), strong chromospheric emission in the H$\\alpha$ Balmer line, flares, and modulations on the light curves caused by the presence of starspots on the stellar surface.  Over the past decades, substantial progress towards the understanding of the structure and evolution of low-mass stars has been made, also including activity effects. \\citet{Chabrier2007} showed that the introduction of rotation and/or magnetic field effects on the models of \\citet{Baraffe1998} could explain the discrepancies between the observed and theoretically predicted mass-radius relationship of EBs. Specifically, they presented two scenarios considering: {\\em (1)} that the effect of magnetic field and rotation alter the efficiency of convective energy transport, which can be modelled by setting the mixing length parameter ($\\alpha$ from Mixing Length Theory) to lower values than those used for solar models; and {\\em (2)} that the stellar magnetic activity present on these objects can be associated with the appearance of dark spots covering the radiative surface, which can be modelled by assuming a new stellar luminosity, $\\hat{L}\\propto \\left( 1 -\\beta \\right) \\hat{R}^{2} \\hat{T}_{\\rm eff}^{4}$, where $\\hat{R}$ and $\\hat{T}_{\\rm eff}$ are the modified stellar radius and effective temperature, respectively, and $\\beta$ is the factor of spots blocking the outgoing luminosity due to their lower temperature (see Section~\\ref{sec:lc} for further details on this parameter).  The results of \\citet{Chabrier2007} show that both these scenarios predict larger radius than standard stellar models, but while the effect of spots is significant over the entire low-mass domain, the effect of convection is relatively smaller for fully convective stars ($M \\le 0.4 M_\\odot$) because convection is nearly adiabatic and changing the mixing length parameter has a modest impact on their stellar structure. Since radius discrepancies have been reported over a wide range of masses, this may be indicating that either the contribution of spots or a combination of spots and activity effect on convective transport are needed to reconcile models with observations.  New accurate fundamental properties of low-mass EBs, and especially the case of CM~Dra \\citep{Morales2009a}, providing precise characterization of two fully convective stars, makes it worthwhile to carry out a new thorough analysis of the agreement between observation and theory using the scenario suggested in \\cite{Chabrier2007}. In Section~\\ref{sec:sample} we briefly describe the discrepancies between stellar structure models and current observations of low-mass stars and we select the sample of best known EB systems to constrain stellar models. In Section~\\ref{sec:lc} we analyze the effect of the photospheric spots on EB light curves and on the parameters fitted to these curves. Critical comparison with models including activity effects is discussed in Section~\\ref{sec:models}, and finally, the conclusions are presented in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  After evaluating different scenarios to explain the discrepancies reported between observation and theory for low mass stars, a rather complex picture is now emerging since no single effect can account for the full size of the differences. Our analysis shows that the $\\sim 5-10$\\% radius discrepancy in M-type EBs is explained by combining three factors: {\\em (1)} A systematic overestimation of the stellar radius in EB light curve analysis caused by the presence of polar spots with the subsequent loss of circular symmetry (amounting to about $\\sim 3$\\%); {\\em (2)} an increase of stellar radii to compensate for the loss of radiative efficiency because of starspots (reproduced by models considering a $\\beta$ parameter of 0.17 that explains 2\\% of the radius difference); and {\\em (3)} an increase of the radius caused by the lower convective efficiency in fast rotators with supposedly strong magnetic fields (reproduced by models through changing the $\\alpha$ parameter) and accounting for 0--4\\% of the radius difference. This scenario also reproduces the $M-\\Teff$ distribution of EBs although effective temperatures are not as fundamentally determined as radii.  Our overall analysis ensures self consistency by comparing with the constraints set by the observations. Predicted spot coverages and simulated light curves are found to reproduce the modulation amplitudes in the real light curves of the EB systems. Also the $\\sim 35$\\% surface coverage is found to be compatible with results from Doppler tomography of active stars. Even relatively inactive systems such as some hosts of transiting exoplanets are found to have significant spot coverages when applying transit mapping procedures \\citep{Pont2007,Silva-Valio2008}.  It is important to note that our results are based on two main hypotheses: {\\em (1)} That {\\em all} of these active EB components have spots that preferentially occupy locations close to the pole (with consequent effect on the light curve analysis) in line with results from Doppler imaging; and {\\em (2)} that {\\em all} EB components have similar surface spottedness caused by their saturated activity regime. The first hypothesis relaxes the classical discrepancy with models to a lower 2--7\\%, and the second hypothesis ensures that $\\beta$ will roughly be ``universal'' for these strongly active systems. In this framework, the system CM Dra is key in reaching the conclusions above. This is because it allows us to separate the effects of the two free parameters, namely $\\beta$ and $\\alpha$, since the components of this system are expected to be fully convective and little affected by changes in the mixing length parameter.  The analysis of low-mass stellar evolution models clearly shows that activity plays an important role on the structure of stars, and that when considering activity effects on stellar structure, models can reproduce the observed values of radii and effective temperatures of low-mass stars. Accurate masses and radii of new low-mass systems coming from missions such as CoRoT and Kepler should improve the statistics and better define the scenario proposed above, especially if comparable active and non-active EB systems are found. Doppler imaging of the EBs analyzed or new methods of mapping stars using interferometry \\citep{Monnier2007} should also give indications on the spot distribution on the stellar surfaces. All these efforts are valuable to understand the vast majority of the stars that populate the Galaxy and, notably, have an impact in the area of exoplanet research, in which planet properties are directly linked to our knowledge of the host star."
    },
    "1005/1005.2100_arXiv.txt": {
        "abstract": "We present results on the evolution of the intrinsic scatter of black hole masses considering different implementations of a model in which black holes only grow via mergers. We demonstrate how merger driven growth affects the correlations between black hole mass and host bulge mass. The simple case of an initially log-normal distributed scatter in black hole and bulge masses combined with random merging within the galaxy population results in a decreasing scatter with merging generation/number as predicted by the Central-limit theorem. In general we find that the decrease in scatter $\\sigma$ is well approximated by $\\sigma_{\\mathrm{merg}}(m) \\sim \\sigma_{\\mathrm{ini}} \\times (m+1)^{-a/2}$ with $a = 0.42$ for a range of mean number of mergers $m < 50$. For a large mean number of mergers ($m > 100$) we find a convergence to $a = 0.61$. This is valid for a wide range of different initial distributions, refill-scenarios or merger mass-ratios.  Growth scenarios based on halo merger trees of a $(100\\ \\mathrm{ Mpc})^3$ dark matter $\\Lambda$CDM-simulation show a similar behaviour with a scatter decrease of $a = 0.30$ with typical number of mergers $m < 50$ consistent with random merging (best matching model: $a=0.34$). Assuming a present day scatter of $0.3$ dex in black hole mass and a mean number of mergers not exceeding $m = 50$ our results imply a scatter of $0.6$ dex at $z = 3$ and thus a possible scenario in which overmassive (and undermassive) black holes at high redshift are a consequence of a larger intrinsic scatter in black hole mass. A simple toy model connecting the growth of black holes to the growth of $\\Lambda$CDM dark matter halos via mergers, neglecting any contribution from accretion, yields a consistent $M_{\\bullet}-M_{Bulge}$ relation at $z=0$ - if we assume the correct initial relation. ",
        "introduction": "There is growing observational evidence that most if not all bulges and elliptical galaxies host a supermassive black hole (SMBH) at the present time (\\citealp{Magorrian}, \\citealp{Genzel}). Furthermore there exists a strong correlation between black hole masses and properties of the host galaxy, e.g. the host bulge mass and the bulge velocity dispersion (\\citealp{Haering}, \\citealp{FerrMerr}, \\citealp{Gebhardt00}, \\citealp{Tremaine02}) and possibly the host halo (\\citealp{Ferr02}) in nearby galaxies. \\citet{Haering} find the relation between black hole mass $M_{\\bullet}$ and bulge mass $M_{\\mathrm{Bulge}}$ to be: \\begin{equation} \\log(M_{\\bullet}/M_{\\odot}) = 8.20 + 1.12 \\times \\log(M_{\\mathrm{Bulge}}/10^{11} M_{\\odot}). \\end{equation} The correlation between black hole mass $M_{\\bullet}$ and velocity dispersion $\\sigma_*$ can be written as \\begin{equation} \\label{BHsigmarel} \\log(M_{\\bullet}/M_{\\odot}) = a  \\log (\\sigma_*/200 \\mathrm{km s^{-1}}) + b, \\end{equation} where a is the slope and b is the zero point. In the literature the values for the slope vary between $a=3.68$ and $a = 4.86$. For example \\citet{FerrMerr} give values for the zero point $b=-2.9$ and the slope $a = 4.80$. \\citet{Gebhardt00} claim a smaller slope of $a=3.75$, \\citet{Tremaine02}  find $a=4.02$ and more recently \\citet{Ferrarese05} estimated a slope of $4.86$ and \\citet{Graham08} one of $3.68$ for barless galaxies. Concerning the intrinsic scatter $\\sigma$ in black-hole mass, $\\log M_{\\bullet}$, most studies agree that the scatter is not larger than $0.3\\ \\mathrm{ dex}$ \\citep{Gebhardt00, Tremaine02, Novak06}. Note, that with $\\sigma_*$ we describe the central velocity dispersion of a galaxy, whereas we characterize the intrinsic scatter as $\\sigma$. In contrast to \\citet{Gebhardt00} and \\citet{Tremaine02}, a recent study by \\citet{Gueltekin09} obtains for Eq. \\ref{BHsigmarel} a slope $a = 4.24$ and an intrinsic scatter in black hole mass of $\\sigma = 0.44\\ \\mathrm{dex}$ for a sample of 49 $M_{\\bullet}$-measurements. For a subsample of early-type galaxies they find a smaller slope as well as a smaller scatter ($a = 3.96, \\sigma = 0.31\\ \\mathrm{dex}$) than for the full sample.  For non-elliptical galaxies the intrinsic scatter is larger with $\\sigma = 0.53\\ \\mathrm{dex}$. In the following the scatter $\\sigma$ will always be given in $\\mathrm{dex}$. The existence of these tight correlations strongly suggest a co-evolution of the black hole and the bulge component of the host galaxy. However, the origin of these relations is uncertain and a subject of current research (e.g. \\citealp{Volonteri09, Pengrandom, Burkert01, Springel05, Johansson09}). Apparently, the origin of these relations can be connected to the gas dynamics in major galaxy mergers (\\citealp{Mihos96}, \\citealp{Naab06a}, \\citealp{Robertson06a}, \\citealp{Hopkins08679}). In this scenario, the black holes grow significantly in gas rich mergers of disk galaxies and the remnants appear on the observed scaling relations. Subsequent, possibly gas poor, major and minor merging (\\citealp{Khochfar03}, \\citealp{Naab06}, \\citealp{Khochfar09}, \\citealp{Naab09}) conserves the relation (\\citealp{Sesana04}, \\citealp{Robertson06a}, \\citealp{Pengrandom}, \\citealp{Hopkins07b}, \\citealp{Hopkins07a}, \\citealp{Springel05}, \\citealp{Johansson09}). \\begin{figure} \\begin{center} \\epsfig{file=Vergleich_Deltalog_onlysigma_Gultekin.eps, width=0.45\\textwidth} \\caption{Comparison of the intrinsic scatter from observations (\\citealp{Tremaine02}, \\citealp{Gueltekin09}) and theoretical predictions (\\citealp{Malbon}). From the 10-90 percentile spread in the model we have calculated the corresponding theoretical scatter. The shaded areas illustrate the error on the observationally estimated intrinsic scatter.} {\\label{Comparison}} \\end{center} \\end{figure}  Observationally the black hole mass relations are well constrained only in the nearby universe and it is unclear if and how they evolve with cosmic time. Several authors have found evidence that galaxies at higher redshift have a higher black hole to bulge mass ratio $M_{\\bullet}/M_{\\mathrm{Bulge}}$ than ellipticals today (\\citealp{McLure06, Treu07, Woo08, Walter, Schramm, Pengobs, Greene09, Natarajan, Salviander07, Shields06}). For a sample of Seyfert galaxies at moderate redshifts $z < 0.1$ the black holes are more massive by $\\Delta \\log M_{\\bullet} \\sim 0.5\\ \\mathrm{dex}$  compared to the local black hole-bulge mass relation (\\citealp{Treu07}, \\citealp{Woo08}). \\citet{Salviander07} find at redshift $z \\approx 1$ an evolution of the $M_{\\bullet}$-$\\sigma_*$-relation by $0.2\\ \\mathrm{dex}$ in black hole mass. At higher redshifts of $z \\sim 2$ \\citet{McLure06} observe black holes $8$ times more massive than expected and \\citet{Pengobs} show that for $z \\ge 2$ the $M_{\\bullet}/M_{\\mathrm{Bulge}}$-ratio is $3-6$ times larger than today. This has been confirmed by \\citet{Greene09} based on a lensed quasar sample. \\citet{Schramm} find evidence for an excess in $M_{\\bullet}/M_{\\mathrm{Bulge}}$ at $z \\sim 3$ of a factor of $\\sim 10$. At redshifts $4 < z < 6$ \\citet{Shields06} obtain for black hole masses in the range of $8 < \\log(M_{\\bullet}/M_{\\odot}) < 10$ a deviation from the present-day $M_{\\bullet}$-$M_{\\mathrm{bulge}}$-relation of $\\Delta \\log(M_{\\bullet}) \\sim 2\\ \\mathrm{dex}$. \\citet{Walter} report an even higher redshift object, a quasar at $z=6$, whose black hole is about $20$ times more massive than expected. Considering a present-day scatter according to \\citet{Tremaine02} ($\\sigma = 0.3$) or according to \\citet{Gueltekin09} ($\\sigma = 0.31$ for ellipticals) all observed black holes at $z \\ge 2$ are outside the $2-\\sigma$ range of the present-day scatter. Furthermore, recent observations (\\citealp{Alexander08, Shapiro09}) show also the existence of under-massive black holes at high redshifts. \\citet{Alexander08} find black hole masses, which are $3$ times smaller than those found in comparable massive galaxies in the local Universe. The results in \\citet{Shapiro09} show that black hole masses at $z=2$ are an order of magnitude lower than those predicted by local scaling relations.  The most obvious explanation for the over-massive black holes at high redshifts are possible selection effects. It is more likely to detect the most luminous and most massive black holes at high redshift than less luminous ones. However, the probability for finding a massive black hole in the mass range $10^9 - 10^{10} M_{\\odot}$ in the observed volume at $z = 3$ (e.g. \\citealp{Schramm}) is extremely low as estimates of the local SMBH mass function from SDSS (\\citealp{Benson07}) would predict no high mass black holes to be found in a similar volume even assuming no evolution in the SMBH mass function. Cosmic variance is very unlikely to be an explanation for the observed, massive black holes. \\citet{Lauer07} point out that there is an additional bias which is due to different selection effects for high-redshift (e.g. black holes in high-z galaxies selected by nuclear activity) and local samples (e.g. black holes in local galaxies selected by luminosity or velocity dispersion). They deduce that because of this bias $M_{\\bullet}$ will typically appear to be too large in a distant sample for a given luminosity or velocity dispersion. Some authors (e.g. \\citealp{Croton06}) explain observed, over-massive black holes at high redshift with a shifted relation towards higher black hole masses for a given bulge mass. He uses the Millennium $\\Lambda$CDM-simulation (\\citealp{Springel01GAD}) coupled with a model of galaxy formation, where galaxy mergers are the primary drivers for black hole and galaxy growth and explore an additional growth channel through which only bulges gain mass, e.g. the disruption of stellar galactic disks in major disk mergers. He argues, that if the bulge growth rate from such disrupted disks is not constant with time, an evolution in the $M_{\\bullet}$-$M_{\\mathrm{bulge}}$-relation can occur. Furthermore, \\citet{Robertson06b} find from simulations of galaxy mergers, that the $M_{\\bullet}$-$\\sigma_*$ -relation shows a slight redshift evolution towards higher black hole masses for a given $\\sigma_*$ at higher redshifts, but they predict no evolution for the $M_{\\bullet}$-$M_{\\mathrm{Bulge}}$-relation. However, \\citet{Hopkins07} using again simulations of major galaxy mergers show that high redshift black holes will be more massive at a fixed bulge mass than expected from the present-day relation. They find an evolution towards lower black hole to bulge mass ratios with cosmic time which is driven by the fact that disks (merger progenitors) have characteristically larger gas fractions at high redshifts. We want to point out that these studies (\\citealp{Croton06, Robertson06b, Hopkins07}) are consistent with each other, they are only focusing on different aspects of the evolution of the relation. However, these methods do not provide a sufficient explanation for the observed under-massive black hole masses at high redshifts, since these methods find to have black holes mainly lying above the median $M_{\\bullet}$-$M_{\\mathrm{Bulge}}$-relation at high redshift evolving towards the relation with decreasing redshift. However, they also predict that the black holes go through a rapid growth phase before they end up above the relation; these objects might be under-massive. Furthermore, if a black hole at high z is already above the median $M_{\\bullet}$-$M_{\\mathrm{Bulge}}$-relation it will not end below the median relation at z=0 if only assuming merging and not gas accretion. Therefore, an alternative explanation for the observed high \\textit{and} low $M_{\\bullet}/M_{\\mathrm{bulge}}$-ratio at high redshifts could be the existence of a larger intrinsic scatter in black hole mass, even assuming no evolution of the mean relation with cosmic time, what would be in agreement with \\citet{Lauer07}.  In this paper we address the question: how does the intrinsic scatter in black hole mass evolve and change with time assuming that black holes grow only via mergers? An answer to this question is also important with respect to similarities and differences between the observed scatter of black hole masses and predictions from theoretical models. \\citet{Malbon}, using semi-analytic modelling, find that the present day scatter in black hole mass decreases significantly with increasing black hole mass. This is in contrast to observations using the full samples of e.g. \\citet{Gebhardt00}, \\citet{Tremaine02}, and \\citet{Gueltekin09}. Here, the scatter appears to be independent of black hole mass. In particular at the high mass end the observed scatter is much larger than the model predictions (see Fig. \\ref{Comparison}). However, \\citet{Gueltekin09} demonstrate that the scatter for non-elliptical galaxies (typically at lower masses) is larger than for elliptical galaxies. This is, at least qualitatively, in agreement with the predictions from \\citet{Malbon}. However, in particular at the high mass end the observed scatter is significantly larger than the model predictions (Fig. \\ref{Comparison}).  So far the time evolution of the scatter in black hole mass has not been investigated in detail. \\citet{Pengrandom} deals with the evolution of the scatter, but he mainly focuses on the aspect of how the present day $M_{\\bullet}$-$M_{\\mathrm{bulge}}$-relation can form in a simple model applying random merging of galaxies. He claims that the relation develops even if black holes and bulges are uncorrelated or incorrectly correlated in the beginning. This behaviour is supposed to result from an initially exponentially decreasing SMBH mass function where minor mergers drive the objects towards the observed correlation. Furthermore, the scatter in black hole mass decreases with increasing merger number, and according to his results the decrease in scatter is dominated by major mergers. However, a quantitative study of the scatter evolution was not presented, which is the main subject of this paper. A further major difference between our work and that of \\citet{Pengrandom} is that we include - besides Monte-Carlo generated random merging scenarios - merging as it is found for dark matter haloes in large scale \\textit{cosmological} simulations.  In the following we will investigate the evolution of the intrinsic scatter assuming: \\begin{itemize} \\item Simple random merging (section \\ref{depletion}) \\item Modified random merging (section \\ref{replenishment}-\\ref{Major_minor_Rep}) \\item Merging in $\\Lambda$CDM-simulations (section \\ref{CDMmerging}) \\end{itemize} We point out that random merging does not describe a full physical evolution process according to currently favoured structure formation models. In principle the model follows dry merging of galaxies and therefore is limited to high mass galaxies, since merging at the high mass end is assumed to be almost dry, so that gas physics and star formation do not play an important role and can be neglected in the growth processes. However, an advantage of using a simple model such as random merging is that we can study separate effects on the scatter evolution, e.g. the influence of the initial mass distribution, of the merger mass-ratio or different refill-scenarios. Then we compare these results to merging according to currently favored structure formation models based on dark matter $\\Lambda$CDM-simulations. Since we know that a significant contribution to black hole growth is caused by accretion, we will discuss this issue in section \\ref{conclusion}. ",
        "conclusions": "In this paper we investigated the evolution of the intrinsic scatter in black hole mass under the assumption that the black holes only grow by mergers with other black holes during galaxy mergers. For many different idealized random merging models with (replenishment) and without (depletion) refilling from an external galaxy pool we find the following general results:  \\begin{enumerate}  \\item The evolution of the black hole distribution can be well described within the framework of the central limit theorem (CLT). Independent of the initial distributions, e.g. log-normal or Schechter, of black holes the distribution always evolves into a normal distribution after a few merger generations. We consider all mergers independent of the mass ratio that galaxies $> 10^{4.7}$ M$_{\\odot}$ experience.  \\item \\label{scatterdecrease} For all random merging models we found a decreasing scatter $\\sigma$ with increasing merging generation $n$ and with increasing merger number $m$. As a consequence of the mass built-up during the merger events the scatter also decreases with increasing black hole mass. Motivated by the CLT, we can approximate the scatter dependence on the mean number of mergers by \\begin{equation} \\sigma_{\\mathrm{merg}}(m) \\approx \\sigma_{\\mathrm{ini}} \\cdot (m+1)^{-a/2}. \\end{equation} Here the exponent $a$ is a measure of the strength of the scatter decerease. For the different random merging models we find $0.30 < a < 0.61$ for a large number of mergers ($m >100$) independent of the initial scatter applied to black hole and bulge masses. In general, replenishment models show a weaker scatter decrease.   \\item Considering either only major or only minor mergers for galaxy growth we found that minor mergers lead to a much stronger scatter decrease than major mergers; hence the smaller the mass ratio of merger events, the more rapidly the scatter decreases. This is in contrast to findings of \\citet{Pengrandom}.  \\item For different replenishment models we found that the higher the refill-ratio and the smaller the typical mass of black holes in the refill-pool, the more slowly the scatter decreases.  \\end{enumerate}  Studying the effect of merging according to current structure formation models in $\\Lambda$CDM-simulations, we find a qualitatively similar behaviour. The scatter decreases with the number of mergers ($a = 0.3$), and as a consequence it also decreases with cosmic time. This is also quantitatively consistent with the best-matching random merging model (refill-ratio $1/3$ and low mass refill-pool) at least for the limit of the low merger number range ($a = 0.34$), since in $\\Lambda$CDM-merging the most massive galaxies experience only $\\sim 50-60$ merger events. Therefore, the scatter evolution in $\\Lambda$CDM-merging can be well approximated by a simple model assuming random merging of galaxies.   From the above results we can draw some implications about recent observations of high redshift black holes: \\begin{itemize} \\item For a simple merger driven growth of black holes we predict that the scatter in black hole mass must have been larger at higher redshift. Assuming an initial scatter of $0.6$ at $z=3$ the over-massive black holes investigated by \\citet{Schramm}, \\citet{Pengobs} and \\citet{McLure06} would be within the $2-\\sigma$ range of this large initial scatter. If these objects had on average $50-60$ dry merger events then we also obtain the present-day scatter value of $\\sim 0.31$ for massive ellipticals \\citep{Gueltekin09}. This shows that the observations of over-massive black holes at high redshifts are consistent with a population of galaxies that has a large scatter in black hole mass ($\\sim 0.6$) at high redshifts being reduced through subsequent merging.  \\item A further important advantage of our model is that we are not only able to explain over-massive black holes at high redshifts, but we can also account for the observed under-massive black holes at $z=2$ (\\citealp{Alexander08, Shapiro09}). These objects would be  lying within the $2-\\sigma$ range of the scatter at $z=2$ ($\\sigma \\approx 0.5$) assuming an initial scatter of $0.6$ at $z=3$. However, we want to point out, that models assuming an evolution of the $M_{\\bullet}-M_{bulge}$ relation (from higher towards lower black hole masses with decreasing redshift) can be an explanation for over-massive black holes at high redshifts, but \\textit{not} for the observed under-massive ones.  \\item We can explain the scatter evolution of model predictions from GALFORM (\\citealp{Malbon}): the decreasing scatter with increasing black hole mass (see Fig. \\ref{Comparison}). This is the consequence of the merger driven growth of black holes. More massive objects have undergone more merger events than less massive ones; this higher merger number leads to a stonger decrease in scatter. However, observations show a larger scatter in black hole masses, especially at the high mass end (\\citealp{Tremaine02}, \\citealp{Gueltekin09}). That might indicate that in the model of \\citet{Malbon} either the number of merger events for massive objects was overestimated (unlikely) or that they started with a too small scatter in black hole mass at high redshift.  \\item Our finding of a decreasing scatter with increasing black hole mass as a consequence of subsequent merging are also consistent with the results of \\citet{Gueltekin09}, who distinguish between two subsamples of ellipticals and non-ellipticals. For ellipticals they find a smaller scatter in black hole masses than for non-ellipticals. Referring elliptical galaxies mainly to the high-mass end and non-elliptical ones mainly to low-mass end, these observations are also consistent with our results. Furthermore, ellipticals presumably had more mergers.  \\end{itemize}   The aim of this paper is to investigate the scatter evolution of supermassive black hole for dry merger driven growth. Despite the interesting insights we have neglected another important physical growth mechanism, e.g. the accretion of gas onto black holes. The importance of an additional growth mechanism is also confirmed by the comparison of the black hole mass function resulting from merging only with observations at $z=0$: we under-estimate the masses of black holes. Furthermore, if we consider only active galaxies that have undergone a merger event in the last $100 \\cdot 10^6\\ \\mathrm{yrs}$ with a merger ratio smaller than $1:10$, we again deviate from observations, which could be a consequence of underpredicting the over-all black hole mass function. Alternatively, it might indicate that merger events are not frequent enough to be an explanation for the observed active galaxies. This implies that there has to be additional gas accretion even if no merger event happened.  Furthermore, according to the Soltan argument (\\citealp{Soltan}) accretion on to massive black holes is the dominant source of energy produced by quasars ($L \\sim \\dot{M_{\\bullet}} c^2$), which presumably indicates that the mass in black holes is mainly generated by accretion of gas. This gas accretion might be triggered by galaxy mergers and their ability to drive large amounts of gas to the centers of the galaxies and possibly feed a black hole (\\citealp{Springel05, DiMatteo05, Johansson09, Hopkins08a}). It was shown in hydrodynamical galaxy-galaxy merger simulations with self-regulated black hole growth that after the active phase of the quasar, star formation and black hole growth are quenched by gas heating through energy release by the active quasar. Following the merger, the quasar host becomes a red and dead, massive elliptical galaxy without any further black hole growth through gas accretion (\\citealp{Johansson09}). For these galaxies the further evolution of their black holes is mainly dominated by dry merger events. However, gas-rich mergers and cold accreation flows seem to be the dominant growth mode for massive high redshift black holes, as the galaxies are more gas rich (\\citealp{Khochfar06, Khochfar09}). The quasar activity peaks at $z\\approx 2-3$ (e.g. \\citealp{Hasinger06, Ueda05}). Therefore, this phase is most likely to be responsible for the scatter in black holes at this redshift (corresponding to the initial scatter in our models). According to \\citet{Hasinger06} the emissivity of the high luminosity objects ($\\log({L/L_{\\odot}}) \\ge 45$) drops by $\\sim 10^2\\ \\mathrm{erg s}^{-1} \\mathrm{Mpc}^{-3}$ between $z \\sim 2$ and $z \\sim 0$. Therefore, the subsequent evolution, at least for these massive black holes, might be driven mainly by dry mergers as investigated in this paper. In addition, there is theoretical evidence that for $z \\le 2$ gas is more efficiently heated by accretion shocks associated with gravitational heating in massive ($M_{\\mathrm{DM}} \\ge 10^{12} M_{\\odot}$) galaxy halos (\\citealp{Khochfar08, Dekel08, Naab09_2, Birnboim07}), suppressing further gas cooling and star formation as well as accretion onto the central supermassive black hole. Thus, 'red and dead' massive spheroids evolve starting at $z \\sim 1$, where again merging may dominate the further growth of these black holes (\\citealp{Dekel06, Khochfar09}). From disk galaxy merger simulations in \\citet{Johansson09_2}, we know that overmassive black holes lying above the black hole-bulge relation do not evolve onto the relation considering only gas accretion driven growth. A solution for this problem could be that the evolution of overmassive black holes onto the relation might be caused mainly by growth through merger events leading to a scatter decrease in black hole mass as shown in this paper. That would also indicate, that gas accretion does not play the most important role in the evolution process of overmassive black holes, at least at late stages, where dry merging is the dominating process.  Intentionally, in this paper we kept the physics very simple and focused only on the growth through merging in order to understand the influence of {\\textit{merging}} on the evolution of the scatter in black hole mass. However, gas accretion is - even for $z<2$ - an important growth channel. The impact of gas accretion on the evolution of the scatter in the $M_{\\bullet}-M_{bulge}$ relation is complicated and so far not easy to assess. To get an idea for such an influence, we refer the reader to \\citet{Johansson09_2}. Using merger simulations of disk galaxies with gas ($3\\cdot 10^9 M_{\\odot} < M_{\\mathrm{bulge}} < 10^{10} M_{\\odot}$) they show two possible scenarios: Starting with galaxies on the relation leads to an increase of the scatter which gets even larger for higher gas fractions. However, considering initially galaxies below and above the relation causes a decrease of the scatter. The decrease seems to get stronger for lower gas fractions. In order to really understand the effect of gas physics we definitely need further investigations of this process in a statistical sense."
    },
    "1005/1005.2624.txt": {
        "abstract": "%***Dick: I did not what to run Latex, so I have just made suggestions by comments in the tex file.  We investigate the coupling between rock-size solids and gas during the formation of gas giant planets by disk fragmentation in the outer regions of massive disks. In this study, we use three-dimensional radiative hydrodynamics simulations and model solids as a spatial distribution of particles.   We assume that half of the total solid fraction is in small grains and half in large solids.  The former are perfectly entrained with the gas and set the opacity in the disk, while the latter are allowed to respond to gas drag forces, with the back reaction on the gas taken into account. To explore the maximum effects of gas-solid interactions, we first consider 10cm-size particles.  We then compare these results to a simulation with 1 km-size particles, which explores the low-drag regime. We show that (1) disk instability planets have the potential to form large cores due to aerodynamic capturing of rock-size solids in spiral arms before fragmentation; (2) that temporary clumps can concentrate tens of $M_{\\oplus}$ of solids in very localized regions before clump disruption; (3) that the formation of permanent clumps, even in the outer disk, is dependent on the grain-size distribution, i.e., the opacity; (4) that nonaxisymmetric structure in the disk can create disk regions that have a solids-to-gas ratio greater than unity; (5) that the solid distribution may affect the fragmentation process; (6) that proto-gas giants and proto-brown dwarfs can start as differentiated objects prior to the H$_2$ collapse phase;  (7) that spiral arms in a gravitationally unstable disk are able to stop the inward drift of rock-size solids, even redistributing them to larger radii; and, (8) that large solids can form spiral arms that are offset from the gaseous spiral arms.  We conclude that planet embryo formation can be strongly affected by the growth of solids during the earliest stages of disk accretion. ",
        "introduction": "Core accretion (e.g., Pollack et al.~1996; Hubickyj et al.~2005) and direct formation by disk instability (e.g, Cameron 1978; Boss 1997, 1998) are both viable modes of gas giant planet formation, and both may be required to explain the diversity of observed planetary systems.    For disk instability to lead to fragmentation, the Toomre (1964) $Q$ parameter must be rapidly lowered in the outer disk, which is where the disk is most likely to have both a low $Q$ and short cooling times relative to the local orbital time (e.g., Rafikov 2007, 2009; Nero \\& Bjorkman 2009; see Durisen et al.~2007 for a review). Although disk instability is quite sensitive to the details of radiative transfer approximations and to opacity (e.g., Johnson \\& Gammie 2003; Cai et al.~2006, 2008, 2009; Boley et al.~2007b;  Mayer et al.~2007; Cossins et al.~2009), hydrodynamics simulations of extended, outer disks demonstrate that fragmentation is possible even under nonisothermal conditions (e.g., Voroybov \\& Basu 2006, 2010; Stamatellos et al.~2007; Stamatellos \\& Whitworth.~2009; Stamatellos et al.~2009; Boley 2009; Forgan et al.~2009; Hayfield  et al.~2010).   If fragmentation is to occur, it is likely to happen during the earliest stages of disk evolution, when mass infall onto the disk can proceed faster than slow, diffusion-like mass transport  (e.g., Boley 2009; hereafter B2009).  This mass loading will likely lead to the development of global spiral modes, nonlocal mass transport, and in some cases, fragmentation. This led B2009 to suggest two modes of gas giant planet formation, with disk instability representing the early mode, operating in newly-formed systems, and with core accretion representing the dominant mode during the evolved stages of disk evolution.  An example of a system with disk instability planets may very well be HR8799 (Marois et al.~2008), as other formation scenarios seem to be unlikely (Dodson-Robinson et al.~2009).  Several observational signatures have been proposed to test whether disk instability plays a role in planet formation. B2009 suggested that a bimodal population of gas giants would reflect that both formation channels are possible and that the ratio of wide orbit gas giants to the total number of planets in a system should increase with decreasing metallicity.  The latter expectation reflects that core accretion should be more sensitive to changes in metallicity than disk instability.  If, instead, wide orbit gas giants are produced mainly by scattering, Dodson-Robinson et al.~(2009) pointed out that  the frequency of systems with gas giants on wide orbits should be dependent on system age.  As more systems with wide orbit gas giants are discovered, these predictions will be tested against observations.  In addition to population trends, metallicity and core size may be a way to distinguish between formation mechanisms. For example,  Helled \\& Bodenheimer (2010) use a planetary evolution code to calculate contraction times for isolated clumps, and follow planetesimal capture of large bodies ($\\ge1$ km) until molecular hydrogen dissociates in the clump's core. This dissociation causes the clump to collapse rapidly to about a Juptier radius, which significantly reduces the cross section of the young planet and marginalizes planetesimal capture.  They find that planetesimal capture can only contribute negligible to modest enhancement of solids for planets at large radii, concluding that planets that form on wide orbits by disk instability should have nearly the same composition as their host star.  If this is correct, metallicity could be a valuable tool in determining formation modes for wide orbit gas giants, as gas giant planets that form by core accretion are expected to be metal rich (e.g., Pollack et al.~1996).  However, a limitation of the Helled \\& Bodenheimer calculation is that their simulations are not within the context of a gravitationally unstable disk.  When clumps form, they tend to form near corotation or at intersections of spiral arms (Durisen et al.~2008), yielding initial fragment masses that contain material from a small segment of the spiral arm (Boley et al.~2010).  Because  spiral arms are the most likely regions for fragmentation, one must also consider enrichment of solids at birth.  Solids are expected to move relative to the gas due to the lack of pressure support (Weidenschilling 1977).  Depending on the direction of the pressure gradient, this disparity between gas and particle forces can lead to either  a tail or headwind for a given particle. The wind exchanges angular momentum with the solid and causes it to migrate.   Haghighipour \\& Boss (2003) showed that this effect will cause particles to become trapped in local pressure maxima, and Rice et al.~(2004, 2006) used SPH calculations with solid particles to demonstrate that spiral arms in a gravitationally unstable disk are ideal places for trapping solids.  These results suggest that enrichment at birth could provide newly-formed fragments with a supernebular solids-to-gas ratio.  In addition,  such concentration of solids could provide enough material to form substantial rocky/icy cores, regardless of whether the fragment survives in the disk.  In this paper, we present hydrodynamics+particle simulations to explore the heavy-element enrichment  of fragments at birth in spatially extended, massive disks during the accretion phase. In Section 2, we outline the numerical methods for this study, including introducing our algorithm for evolving solids.  We then describe the numerical experiments and disk initial conditions in Section 3.  We present our results in Sections 4 and 5, and discuss the implications of some of these results in section 6.  We summarize our conclusions with a bulleted list in Section 7. ",
        "conclusions": "We have investigated the coupling between rocks and gas in fragmenting, gravitationally unstable disks.  We have explored the behavior of 10 cm-size and 1 km-size particles in this study.  The former corresponds to sizes that are expected to be swept up efficiently by spiral arms, while the latter is largely decoupled from gas drag.  Because the small-grain limit is simply entrained with the gas, these simulations explore the two nontrivial regimes for the behavior of solids in gravitationally unstable disks.  Our simulations have many implications for disk fragmentation and planet formation, which we summarize in bullet form.  \\begin{itemize}  \\item Spiral arms concentrate rock-size material into the regions that are most likely to fragment.  However, this aerodynamic capturing of solids leads to, at best, a total enrichment of a clump by a factor of two.  A distribution of solids that does not contain a large amount of rock-size particles will reduce this effect, without considering solid evolution, as demonstrated by the 1 km-size run. \\item Subsequent mass growth does not dilute the birth enrichment in these simulations for the 10 cm-size solids because clumps accrete gas from filaments, which are rock enriched, or grow by collisions with spiral arms and/or other clumps, which are also rock enriched. \\item If most of the solid mass is in 1 km-size solids or in sizes that are equally decoupled from the gas, clumps can become anemic in heavy elements. \\item Rock densities can become comparable to or exceed local gas densities in some structures. Even for the 1 km-size case, regions of supernebular solids-to-gas occur due to differences between the spiral structures in the solids and the gas. \\item The rock distribution appears to affect the behavior of gravitational instabilities and fragmentation.  The influence of the solids on the gaseous disk, particularly fragmentation, may mainly be due to gravitational potential changes.  Additional simulations with distributions of solids need to be explored, but must be coupled with a solid size evolution model. \\item Gravitational instabilities interfere with the simple inward drift of solids.  Instead, rocks that would normally have a very high inward velocity are entrained with the evolution of the spiral arms, and are even transported out to large radii. \\item The fragmentation process is affected by the grain size distribution, i.e., the opacity. For the temperatures in the outer disk, larger grain sizes can stop dense condensations from forming long-lived clumps. \\item Although the total enrichment is modest, solids are highly concentrated near the center of the clump as part of the fragmentation process.  The solids-to-gas ratio can be as high as a few tenths in these regions, with tens of Earth masses of rock available for immediate core formation, i.e., the material does not need to settle from the entire clump.  This could significantly decrease the core formation timescales calculated by Helled \\& Schubert.~(2008).  Overall, the results suggest that disk instability planets can have massive cores. \\item Fragments can become sheared away by tides, but because rocks are concentrated near the center of the clump, these fragments may still form rocky cores before the gas is completely disrupted. If such solid cores survive tidal disruption of the extended gas envelope, they may very well represent the first rocky/icy planets  or planetary embryos in a planetary system.  \\end{itemize}  Understanding the behavior of solids in gaseous disks is required for building a basic picture of planet formation, even in the context of a gravitationally unstable disk.  As demonstrated here, fragments, although only modestly enriched, can be differentiated at birth. This leaves open the possibility of a hybrid scenario that we call core assist plus gas capture.  In addition, it also suggests that clumps as massive as brown dwarfs could also start as differentiated objects, depending on the degree of the enhancement of solids within the spiral arms at the time of gas disk fragmentation.   These formation channels should have observational consequences, which we only begin to touch upon here, but include metallicity and core size, as well as orbital locations if done in a statistical sense. Because the fragmentation process is strongly affected by dust characteristics, future studies should include  grain growth models that self-consistently change dust opacities and capture, even if heuristically, the formation of solids with a distribution of sizes.   Although challenging, such a model will greatly aid in determining the likelihood of fragmentation and what type of objects disk instability will typically form.  ACB and RHD would like to acknowledge the following individuals who have influenced our thinking and have provided us with ideas regarding the topics discussed in this manuscript: F.~Adams, P.~Bodenheimer, A.~Boss, C.~Dullemond,    E.~Ford,  L.~Fouchet, N.~Haghighipour,  L.~Hartmann,   T.~Hayfield,  R.~Helled, Th.~Henning, O.~Hubickyj,  A.~Johansen, H.~Klahr, J.~Lissauer, G.~Lodato,   L.~Mayer, S.~Michael,  F.~Miniati,  G.~Morfill, M.~Payne, K.~Rice,   and D.~Stamatellos.  We would particularly like to acknowledge the hospitality of the Max Planck Institute for Astronomy and %***dick: MPE not MPIE. Actually, I think you should write out the names in their entirety -- \"of the Max Planck Institutes for Astronomy and for Extraterrestrial Physics\". the Max Planck Institute for Extraterrestrial Physics over the years.  ACB would like to thank the Kavli Institute for its hospitality during the ``Theory and Observations of Exoplanets''  workshop.   ACB's support was provided in part by a Theory Fellowship through the University of Florida and in part under contract with the California Institute of Technology (Caltech) funded by NASA through the Sagan Fellowship Program. RHD was supported by NASA Origins of Solar Systems grant NNX08AK36G. This research was supported in part by the National Science Foundation under Grant No.~PHY05-51164. Resources supporting this work were provided by the NASA High-End Computing (HEC) Program through the NASA Advanced Supercomputing (NAS) Division at Ames Research Center.  \\appendix"
    },
    "1005/1005.1932_arXiv.txt": {
        "abstract": "We present extensive radio and X-ray observations of the nearby Type Ic SN\\,2007gr in NGC\\,1058 obtained with the Very Large Array and the Chandra X-ray Observatory and spanning 5 to 150 days after explosion.  Through our detailed modeling of these data, we estimate the properties of the blastwave and the circumstellar environment.  We find evidence for a freely-expanding and non-relativistic explosion with an average blastwave velocity, $\\overline{v}\\approx 0.2c$, and a total internal energy for the radio emitting material of $E\\approx 2\\times 10^{46}$ erg assuming equipartition of energy between electrons and magnetic fields ($\\epsilon_e=\\epsilon_B=0.1$).  The temporal and spectral evolution of the radio emission points to a stellar wind-blown environment shaped by a steady progenitor mass loss rate of $\\dot{M}\\approx 6\\times 10^{-7}~\\rm M_{\\odot}~yr^{-1}$ (wind velocity, $v_w=10^3~\\rm km~s^{-1}$). These parameters are fully consistent with those inferred for other SNe Ibc and are in line with the expectations for an ordinary, homologous SN explosion.  Our results are at odds with those of \\citet{par10} who recently reported evidence for a relativistic blastwave in SN\\,2007gr based on their claim that the radio emission was resolved away in a low signal-to-noise Very Long Baseline Interferometry (VLBI) observation.  Here we show that the exotic physical scenarios required to explain the claimed relativistic velocity -- extreme departures from equipartition and/or a highly collimated outflow -- are excluded by our detailed Very Large Array radio observations.  Moreover, we present an independent analysis of the VLBI data and propose that a modest loss of phase coherence provides a more natural explanation for the apparent flux density loss which is evident on both short and long baselines.  We conclude that SN\\,2007gr is an ordinary Type Ibc supernova. ",
        "introduction": "\\label{sec:intro}  Nearly 25 years have elapsed since Type Ibc supernovae (SNe Ibc) were first recognized as a distinct class of cosmic explosions \\citep{emn+85,fs85,wl85}.  They are now commonly understood to mark the gravitational core-collapse of massive stars from which an explosion launches ejecta to bulk velocities of $\\sim 10,000~\\rm km~s^{-1}$ (e.g.,~\\citealt{fil97}).  Motivated by the observed lack of hydrogen (and often helium) in their optical spectra, the favored progenitors of SNe Ibc are Wolf-Rayet stars that have ejected their massive envelope through strong stellar winds \\citep{bs86} or interaction with a binary companion \\citep{pjh92}. Recently, the observational realization that most long-duration gamma-ray bursts (LGRBs) are accompanied by SNe Ibc has fueled a new surge of interest in these massive star explosions (see \\citealt{wb06} and references therein).  The LGRB-SN connection and the relative rates of these events implies that LGRBs are a rare subclass of SNe Ibc, distinguished by the production of a relativistic outflow that decouples from the bulk SN ejecta during explosion.  For many LGRBs, the energy associated with the relativistic blastwave, $E\\approx 10^{48}-10^{52}$ erg, is comparable to that of the non-relativistic SN ejecta, $E_{\\rm SN}\\approx 10^{51}-10^{52}$ erg (e.g.,~GRB\\,030329/SN2003dh; \\citealt{bkp+03,hsm+03,mdt+03,fsk+05}).  This duality challenges the standard SN mechanism in which a homologous explosion couples at most $0.01\\%$ of the total energy to mildly-relativistic ejecta \\citep{tmm01}.  It is therefore generally accepted that a ``central engine\" --- an accreting and rapidly-rotating compact object \\citep{mwh01} or a magnetar (e.g.,~\\citealt{tcq04}) --- is additionally required to power the energetic and relativistic blastwaves of LGRBs.  The critical question is whether ordinary SNe Ibc also harbor central engines, and in turn, the ability to produce even weak relativistic outflows.  Identifying engine-driven, relativistic explosions requires direct measurements of the blastwave velocity and energy.  While optical SN emission predominantly traces the Nickel-56 synthesized in the bulk SN ejecta \\citep{arn82}, radio and X-ray observations directly probe the synchrotron radiation produced as the blastwave shocks material in the circumstellar medium (CSM;~\\citealt{c82}).  Over the past decade, dedicated radio studies of SNe Ibc have consistently pointed to blastwave velocities of just $\\langle v\\rangle\\approx 0.15c$ and associated energies, $\\langle E\\rangle\\approx 10^{47}$ erg, in agreement with the expectations for an ordinary core-collapse explosion \\citep{c98,bkc02,bkf+03,skb+05,sck+06,cf06, ams07,sbp+08}.  Recently, luminous radio emission was detected from the broad-lined Type Ic SN\\,2009bb pointing to an energetic ($E\\gtrsim 10^{49}$ erg) and relativistic ($v\\gtrsim 0.85c$) outflow powered by a central engine \\citep{scp+10}.  This discovery marked the first relativistic SN explosion identified without a detected gamma-ray counterpart.  Together with the results of on-going radio supernova studies, this result independently confirmed that the fraction of SNe Ibc with engine-driven relativistic outflows is exceedingly small, $\\sim 1\\%$ \\citep{bkf+03,snb+06,scp+10}.  Against this backdrop, \\citet{par10} recently claimed evidence for a mildly-relativistic blastwave ($v > 0.6c$) from the otherwise ordinary Type Ic SN\\,2007gr based on a low signal-to-noise Very Long Baseline Interferometry (VLBI) radio observation.  Here we present an independent analysis of these VLBI data and report in detail on our extensive multi-frequency radio and X-ray observations of SN\\,2007gr obtained with the Very Large Array\\footnote{The Very Large Array and Very Long Baseline Array are operated by the National Radio Astronomy Observatory, a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}  (VLA) and the Chandra X-ray Observatory (CXO).  In contrast to the hypothesis of \\citet{par10}, our multi-frequency radio data ($\\nu_{\\rm obs}=1.4-22$ GHz) and early X-ray observation reveal that the SN emission is best explained by an ordinary, non-relativistic blastwave with an average velocity, $\\overline{v}\\approx 0.2c$, and an associated energy, $E\\approx 2\\times 10^{46}$ erg.  This result stems from the fact that we directly measure the properties of the radio spectral peak across multiple epochs.  We further demonstrate that a conservative interpretation of the VLBI data is fully consistent with non-relativistic SN expansion.  Moreover, we argue that the exotic physical conditions required by a mildly relativistic outflow are ruled out by our VLA observations of the temporal and spectral evolution of the radio emission.  Overall, we find that the SN\\,2007gr blastwave properties are analogous to those of other Type Ibc supernovae and starkly dissimilar from the relativistic jets of LGRBs. ",
        "conclusions": "\\label{sec:conc}  In this paper, we present a critical analysis of blastwave properties of Type Ic SN\\,2007gr following the recent claim by \\citet{par10} that this otherwise ordinary SN produced a relativistic outflow similar to those of nearby LGRBs.  We show that the full dataset of radio (VLA {\\it and} VLBI) and X-ray observations for SN\\,2007gr are more naturally explained by an ordinary, non-relativistic, and homologous SN explosion.  Our conclusions stem from the fact that we directly measured the frequency and flux density of $\\nu_p$ across several epochs.  This result underscores the necessity of multi-frequency and long-term radio monitoring of SNe Ibc in the search for relativistic outflows.  We conclude with the following points:  \\noindent \\begin{enumerate} \\item{A freely-expanding and non-relativistic blastwave model best reproduces our extensive radio and X-ray observations, indicating an expansion velocity of $\\overline{v}\\approx 0.2c$ and energy of $E\\approx 2\\times 10^{46}$ erg.  These blastwave parameters are consistent with those inferred from the bulk ejecta properties given the expected energy profile of the ejecta.}  \\noindent \\item{\\citet{par10}'s proposed mildly-relativistic blastwave velocity would require exotic physical scenarios (severe departures from equipartition or a decelerated, detached blastwave) that are implausible given the requirements of our VLA observations.}  \\smallskip  \\noindent \\item{Through our independent analysis of the VLBI data sets, we confirm the weak detection of a source consistent with our VLA position for SN\\,2007gr.  However, in our conservative interpretation of the low signal-to-noise second epoch detection, we attribute the apparent loss in flux density to modest phase decoherence, instead of relativistic SN expansion.  We also suggest that the WSRT observation is contaminated by diffuse emission from the host galaxy.}  \\end{enumerate}  Finally, we note that while the data for SN\\,2007gr do not point to a relativistic explosion, they do offer new insight on the non-thermal properties of Type Ibc SNe.  In addition to being one of the {\\it nearest} of such explosions discovered to date, SN\\,2007gr is also one of the {\\it least} radio luminous.  This can be directly attributed to its low density circumstellar environment that was shaped by a steady mass loss rate, $\\dot{M}\\approx 6\\times 10^{-7}~\\rm M_{\\odot}~yr^{-1}$ (\\S\\ref{sec:dyn}).  For comparison, this is $10^3-10^4$ times lower than the mass loss rates inferred for radio SNe Ibc with the strongest circumstellar interaction (e.g.,~SNe 2003L, 2003bg;~\\citealt{skb+05,sck+06}) which are preferentially detected in the radio and X-ray bands thanks to their luminous non-thermal emission.  Since SNe with low mass loss rates generally give rise to weak radio signals, they are only detectable nearby ($d\\lesssim 30$ Mpc) with current cm-band facilities.  This, together with the discovery rate of SNe Ibc within $d\\approx 10$ Mpc (a few each decade), statistically suggests that SN\\,2007gr represents one of the most ordinary SNe Ibc studied to date, characterized by explosion and environmental properties that are typical of the bulk population. With the significant improvement in continuum sensitivity enabled by the Expanded VLA (EVLA; \\citealt{pnj+09}), we will soon be able to detect such ordinary SN Ibc to distances of $\\sim 100$ Mpc which will. in turn, broaden our understanding of these unique cosmic explosions.  \\medskip  The authors especially thank Mark Reid for helpful discussions.  We also thank Dale Frail, Edo Berger, Andrew MacFadyen and Eli Waxman.  AMS is supported by a Hubble fellowship. RAC acknowledges support from NASA grant NNG06GJ33G."
    },
    "1005/1005.0966.txt": {
        "abstract": "% context heading (optional) {Gravitational microlensing is not only a successful tool for discovering distant exoplanets, but it also enables characterization of the lens and source stars involved in the lensing event.} % aims heading (mandatory) {In high magnification events, the lens caustic may cross over the source disk, which allows a determination of the angular size of the source and additionally a measurement of its limb darkening.} % methods heading (mandatory) {When such extended-source effects appear close to maximum magnification, the resulting light curve differs from the characteristic Paczy\\'nski point-source curve. The exact shape of the light curve close to the peak depends on the limb darkening of the source. Dense photometric coverage permits measurement of the respective limb-darkening coefficients.} % results heading (mandatory) {In the case of microlensing event OGLE~2008-BLG-290, the K giant source star reached a peak magnification of about 100. Thirteen different telescopes have covered this event in eight different photometric bands. Subsequent light-curve analysis yielded measurements of linear limb-darkening coefficients of the source in six photometric bands. The best-measured coefficients lead to an estimate of the source effective temperature of about $4700^{+100}_{-200}$~K. However, the photometric estimate from colour-magnitude diagrams favours a cooler temperature of $4200 \\pm 100$~K.} % conclusions heading (optional) {As the limb-darkening measurements, at least in the CTIO/SMARTS2 $\\Vs-$ and $\\Is-$ bands, are among the most accurate obtained, the above disagreement needs to be understood. A solution is proposed, which may apply to previous events where such a discrepancy also appeared.} ",
        "introduction": "\\label{sec:intro} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Astrophysical opportunities for measuring stellar limb darkening are very scarce. Moreover, they are often limited to very specific types of stars. Up to recently, most measurements had been based on analyses of eclipsing-binary light curves \\citep{Popper1984}, requiring large amounts of high-precision photometric data but only rarely yielding useful constraints on the recovered limb-darkening parameters \\citep[e.g.][]{PE1981,Southworth2005}.  Resolving stellar surfaces by interferometry is a promising technique, limited nevertheless at present to nearby giants and supergiants. For most of the sufficiently well-resolved stars it has been shown that a fixed model-atmosphere-based limb-darkening profile agreed with the observations better than a uniform or fully limb-darkened stellar disk. Only a limited number of cases have been used to actually measure limb darkening from interferometric data \\citep[e.g.,][]{Burns1997,Perrin2004,Aufdenberg2006,Wittkowski2006}. The same is true for observations of lunar occultation of stars, which yield limb-darkening measurements only rarely \\citep{Richichi1990}.  An alternative method which permits accurate limb-darkening measurement practically independent of stellar type is based on gravitational microlensing. In a stellar gravitational microlensing event the flux from an observed source star is temporarily amplified by the gravitational field of another object passing in the foreground and acting as a gravitational lens \\citep{Paczynski1996}. The lensing object may be typically a single star, a binary star, or a star with one or several planetary companions. A particularly interesting situation occurs in so-called caustic-crossing events, in which the caustic of the lens directly crosses the projected disk of the source star, typically on a timescale of a fraction of a day to a few days. The very high angular resolution provided by the caustic presents a unique opportunity to measure the source star's limb darkening from photometric observations or the centre-to-limb variation of spectral features from spectroscopic observations of the crossing \\citep[e.g.][]{Witt1995,HSL2000,Heyrovsky2003}.  The first well-documented source-resolving caustic-crossing microlensing event was MACHO Alert 95-30 \\citep{Alcock1997}, in which an M4~III giant source star was scanned by the point-like caustic of a single lens. The observations permitted measurement of the source size and detection of its limb darkening. Precise measurement of the limb darkening was not possible due to low-amplitude variability of the source \\citep{Heyrovsky2003}. The first limb-darkening measurement came from the analysis of binary-lens event MACHO~97-BLG-28 \\citep{Albrow1999b}, producing $V-$ and $I-$band coefficients of the linear and square-root limb-darkening laws for the K2~III giant source. By the time of writing, six further binary-lens caustic-crossing events yielded limb-darkening measurements of their source stars: MACHO~97-BLG-41 \\citep{Albrow2000b}, MACHO~98-SMC-1 \\citep{Afonso2000}, OGLE~1999-BUL-23 \\citep{Albrow2001a}, EROS~2000-BLG-5 \\citep{An2002,Fields2003}, MOA~2002-BLG-33 \\citep{Abe2003}, and OGLE~2002-BLG-069 \\citep{Cassan2004,Kubas2005}.  Limb-darkening measurements are also available from four single-lens caustic-crossing events: OGLE~2003-BLG-262 \\citep{Yoo2004}, OGLE~2003-BLG-238 \\citep{Jiang2004}, OGLE~2004-BLG-254 \\citep{Cassan2006,Heyrovsky2008}, and OGLE~2004-BLG-482 \\citep{Zub2009}. In addition, there are a comparable number of events that required the inclusion of limb darkening in their analysis (such as MACHO Alert 95-30), but for which a sufficiently accurate limb-darkening measurement was not (and often could not be) performed. These include two further single-lens events \\citep{Yee2009,Batista2009}, as well as five of the eight published planetary microlensing events, in which a star with a planet acted as the lens \\citep{Bond2004,Dong2009b,Dong2009a,Beaulieu2006,Janczak2010}.  In this paper we report the results of the analysis of OGLE~2008-BLG-290, a microlensing event in which the caustic of a single lens crossed the disk of a K giant in the Galactic Bulge. In \\Sec{sec:data} we describe the obtained data and their reduction. We discuss the properties of the source star in detail in \\Sec{sec:source}. The data modelling in \\Sec{sec:model} is followed in \\Sec{sec:LDmeasure} by a description of the limb-darkening measurements in six photometric bands and their comparison with model atmospheres. We devote \\Sec{sec:lens} to the study of the properties of the lens. We explore the possibility of deviations caused by a potential planetary companion to the lens in \\Sec{sec:planet} and consider other potential systematics in \\Sec{sec:systematics}. The main results are summarized in \\Sec{sec:discussion}.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:discussion} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  We have performed dense photometric monitoring of the microlensing event OGLE~2008--BLG--290, a short duration and very small impact parameter microlensing event generated by a point-like lens transiting a giant star. The peak magnification was about $A(\\to)\\sim 100$, effectively multiplying the diameter of our network telescopes by a factor $\\sim 10$. Using a calibrated colour-magnitude diagram analysis and isochrones, we find a source angular radius $\\ThS \\simeq 6.5 \\pm 0.4 \\,\\muas$, and a physical radius $\\RS \\simeq 10 \\ \\rsun$ at the adopted source distance $\\DS = 7.6$~kpc. A Galaxy model together with the event time-scale $\\tE = 16.4~\\mbox{days}$ and the angular Einstein radius $\\ThE = 0.3$~mas yielded statistically inferred median parameters, namely a lens mass $M\\sim 0.3\\,\\msun$, a lens distance $\\DL \\sim 6.0~\\mbox{kpc}$, and a relative velocity $\\vperp \\sim 190~\\mbox{km}\\,\\mbox{s}^{-1}$.  From our photometric data, we have derived accurate measurements of the linear limb-darkening coefficients of the source for the $\\Vs$, SDSS $r$, $R$, MOA-red, $\\Is$, and $I$ broadband filters. The obtained limb-darkening profiles lead to a source-star temperature estimate of $\\teff=4700^{+100}_{-200}$~K when compared with the limb-darkening profiles of Kurucz's ATLAS9 atmosphere models. This result is in marginal disagreement with the corresponding estimate based on measured colours, a finding that has already been noted for several previous events measured by the microlensing technique, which also parallels the known discrepancies in temperature estimates based on photometry vs. spectroscopy of similar cool giants \\citep{Fulbright2006}. Possible explanations involve reddening corrections which may lead to wrong colour estimates, or inaccuracies in model atmosphere physics.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1005/1005.1874_arXiv.txt": {
        "abstract": "In this paper we consider a quark nova occurring inside an exploding star. The quark nova ejecta will shock when interacting with the stellar envelope. When this shock reaches the surface of the star, the energy is radiated away. We suggest that this energy may be seen in X-rays, and show here that this may explain some flares seen in the X-ray afterglow of long gamma ray bursts (GRBs). A quark nova inside an exploding star need not be followed by a GRB, or the GRB may not be beamed towards us. However, the shock break-out is likely not beamed and could be seen even in the absence of a GRB. We suggest that XRO 080109 is such an event in which a quark nova occurs inside an exploding star. No GRB is formed, but the break out of the shock leads to the XRO. ",
        "introduction": "Quark stars (QSs) are hypothetical objects composed of deconfined u, d, and s quarks. The strange quark matter (SQM) hypothesis states that SQM may be the absolute ground state of strong interacting matter rather than $^{56}Fe$ \\cite{bodmer71, witten84}, and therefore QSs would be stable objects. Even before the formulation of the SQM hypothesis Itoh \\cite{itoh70} discussed the possible existence of QSs.  QS masses and radii are likely comparable to neutron stars (NSs), but slightly smaller radius due to their higher density. A possible formation scenario for QSs is that the central density in a neutron star increases past a critical density at which quarks deconfine. This can happen either due to spin down of a rapidly rotating NS \\cite{staff06}, or through accretion onto the NS. A QS can then be formed in a quark nova (QN) \\cite{ouyed02}, during which up to $10^{53}$ erg can be released in an explosive event. The iron rich NS crust will likely be blown away in the explosion. For a QN to occur soon after the core collapse, either the NS magnetic field must be very high ($B\\sim10^{14}-10^{15}$ G) and the NS rapidly rotating, or more likely through accretion. In \\cite{ouyed09} the interaction of the iron rich QN ejecta and the surrounding star was studied, and applied to long gamma ray bursts (GRBs). There it was suggested that QN ejecta leaving the star through a funnel could lead to precursor activity.   Long GRBs typically have a duration $T_{90}\\sim 100$ seconds, after which the gamma radiation ends. The X-ray telescope on board Swift is sometimes able to focus onto a GRB at a timescale of around a 100 seconds. A generic X-ray light curve \\cite{zhang06} can be constructed based on XRT observations. It shows a sharp drop in the X-ray light curve lasting for a few hundred seconds, followed by a plateau phase extending to several times $10^4$ seconds, before a ``normal'' power law decay is seen. Overlayed on all this are one or more flares, some of which can have a fluence comparable to that of the prompt gamma emission itself yet typically they have a much lower energy. Other bursts do not show the sharp drop and plateau phase, but rather a ``normal'' power law from early times. There may however be flares overlayed also on this light curve.  The ``normal'' power law decay is what is expected from the external shock. This power should be steeper than -0.75 for pre-jet-break and steeper than -1.5 for post-jet-break light curves \\cite{liang08}, yet not steeper than -3 \\cite{troja07}. Hence the external shock alone has difficulty explaining the very sharp drop, the flares, and the plateau phase, which is why other mechanisms involving the inner engine must be sought in the cases where these features are seen.  In \\cite{staff07} a three stage model for the inner engine of long GRBs was suggested. Stage 1 is a NS formed by the collapse of the iron core in a SN. Accretion or spin-down can make the NS explode in a QN leaving a QS behind, which leads to stage 2 that is a jet launched from hyperaccretion onto a QS giving the GRB prompt emission. When the QS has accreted sufficiently, it collapses to a black hole (BH) whereby the third stage can be reached, which is accretion onto a BH. This may launch another jet that can give rise to one big flare (or several flare in rapid succession). If accretion ends before the QS collapses to a BH, stage 3 will never be reached and instead spin down of the highly magnetized rapidly rotating QS may power a secondary outflow giving rise to a flat segment in the X-ray afterglow \\cite{staff08}.  In this paper we build on the work in \\cite{ouyed09} and explore the breakout of the shock formed when the QN ejecta interacts with the stellar envelope. We will use this to explain features seen in GRBs or XRO 080109. The expressions for describing the QN induced shock's propagation through the stellar envelope is given in section \\ref{theorysection}. In section \\ref{grbsection} we study possible observational consequences of the QN that formed the QS leading to the GRB. Using the simple expressions presented in \\ref{theorysection} we show that some later flares in GRBs may be formed by the ejecta launched from the QN interacting with the stellar envelope. The QN ejecta are likely ejected isotropically, so even when no GRB is seen these weaker and later flares can be seen. A QN does not have to be followed by a GRB and in section \\ref{xrosection} we suggest that XRO 080109 was a case where a QN occurred inside a SN, but no GRB was produced. A summary of the three stage model and its features is given in section \\ref{summarysection}. ",
        "conclusions": "\\label{summarysection}  In this paper we have explored the break-out of the shock created by a QN ejecta when it interacts with the stellar envelope surrounding the QN. The QN is assumed to occur inside an exploding star. In the event of a non-uniform envelope the resulting light curve from this shock break-out can become complicated. Here we simplified the non-uniformity to assuming that there are two envelopes, one outside the other, and that a fraction of the QN ejecta interacts with the inner envelope and the rest with the outer. Based on this we find that two flares can result from the shock break-out. We suggest that these flares can explain some of the flares seen in GRB afterglows. The shock break-out is likely not beamed, whereas GRBs are. A GRB not beamed towards us will not be seen, while the QN shock break-out can still be seen. A QN can also occur inside a SN without forming a GRB at all, while the shock break out can be seen. We suggest that this is the case in XRO 080109.  We here summarizes the three stage model, and emphasizes the new aspects proposed in this paper:  \\begin{itemize}  \\item The iron core in an initially massive, rapidly rotating star collapses, creating a (proto-)neutron star. This is the first stage in the model.  \\item Through spin down of the NS or accretion onto the NS (or both), the central density in the QS reaches a critical value at which a QN occurs, creating a highly magnetized QS.  \\item A hyperaccretion disk surrounding the QS accretes onto the QS, launching an ultrarelativistic jet in which internal shocks produces the GRB. This is the second stage in the model.  \\item The jet launching continues until the accretion is no longer able to heat the QS above 7.7 MeV, until accretion is no longer channeled to the polar cap (because the QS magnetic field drops or the accretion rate increases), or until the accretion halts (because there is no disk left). This leads to a sharp drop off due to curvature radiation frequently seen in the X-ray afterglow. If the external shock has already formed and is strong before the sharp drop off, the sharp drop off may not be seen.  \\item Flares seen before the sharp drop off may be caused by the same mechanism that created the GRB itself (internal shocks), only that the radiation is softer.  \\item If sufficient mass is accreted onto the QS, it will collapse to a BH and Stage 3 is reached which is continued accretion onto the BH. This can launch another jet, leading to a major X-ray flare. This flare will likely occur after the sharp drop off started, but not necessarily before it ended (it may also occur after the end of the sharp drop off).  \\item We assume that a maximum of $0.1 M_\\odot$ can be accreted onto the polar cap. Usually accretion rates of the order $\\dot{m} \\gtrsim 10^{-4} M_\\odot/sec$ is needed to explain the observed GRB power, indicating a total duration of the QS phase $t\\lesssim 1000 sec$. However, if the accretion rate drops during stage 2, this duration may be extended.  \\item In the event that the QS survived the accretion stage, stage 3 will not be reached. Because of the accretion, this QS is likely rapidly rotating, and we assume it to be highly magnetized. Hence rapid spin down due to magnetic braking will occur. This can lead to an observed flat segment in the afterglow, if the spin down power outshines the external shock. The QS may reach an unstable state during spin down, leading it to collapse to a BH. The observed light curve will then drop sharply down to the level given by the external shock.  \\item There may therefore be two separate outflows leading to two different light curves seen for instance as the X-ray and optical afterglow. One is created by spin down and the other by the external shock.  \\item In this paper we suggested that the shocks created by the chunks ejected in the QN can create observable flaring when it breaks out. In some cases this may be observed, in others it may be dwarfed by the external shock. The time at which these shocks break out depends largely on the size of the exploding star's envelope, which in turn depends on the delay between the collapse of the iron core and the QN. Here we have showed how two flares may result if the exploding star's envelope is not uniform.  \\item A flatter than expected light curve is also possible with the refreshed shock mechanism, in which slower part of the QS jet catches up with the external shock at later time and ``refreshes'' it leading to a shallower light curve. At the end of the refreshing, a break in the light curve is observed as the light curve steepens when no more material can refresh the external shock.  \\item A jet break can be observed in the light curve when $1/\\Gamma>\\theta_{\\rm jet}$ ($\\Gamma$ being the Lorentz factor of the jet and $\\theta_{\\rm jet}$ being the opening angle of the jet itself) at which point the light curve steepens. Since there can be two separate outflows leading to the optical and X-ray afterglow, this may explain why the breaks (thought to be jet-breaks) are not achromatic.  \\item The QS jet, BH jet, and QS spin down outflow are all assumed to be beamed in two bipolar jets. The QN ejecta on the other hand is assumed to not be beamed, and it could therefore still be possible to see this even when the GRB is not beamed towards us. A QN may also occur inside an exploding star without forming a GRB (this happens if no hyperaccretion disk forms around the QS). However, the chunk break-out may still be seen. We have here suggested that XRO 080109 is such an event in which a QN occurred but no GRB was formed.  \\end{itemize}  \\bigskip The authors would like to thank the organizers of CSQCDII for a very interesting and fruitful meeting. This work has been supported, in part, by grants AST-0708551, PHY-0653369, and PHY-0326311 from the U.S. National Science Foundation and, in part, by grant NNX07AG84G from NASA\u2019s ATP program."
    },
    "1005/1005.1273_arXiv.txt": {
        "abstract": "The Vela-D region, according to the nomenclature given by \\citet{mur:may}, of the star forming complex known as the Vela Molecular Ridge (VMR), has been recently analyzed in details by \\citet{olmi2009}, who studied the physical properties of 141 pre- and proto-stellar cold dust cores, detected by the ``Balloon-borne Large-Aperture Submillimeter Telescope'' (BLAST) during a much larger (55~deg$^2$) Galactic Plane survey encompassing the whole VMR. This survey's primary goal was to identify the coldest, dense dust cores possibly associated with the earliest phases of star formation. In this work, the dynamical state of the Vela-D cores is analyzed. Comparison to dynamical masses of a sub-sample of the Vela-D cores estimated from the $^{13}$CO survey of \\citet{elia2007}, is complicated by the fact that the $^{13}$CO linewidths are likely to trace the lower density intercore material, in addition to the dense gas associated with the compact cores observed by BLAST. In fact, the total internal pressure of these cores, if estimated using the $^{13}$CO linewidths, appears to be higher than the cloud ambient pressure. If this were the case, then self-gravity and surface pressure would be insufficient to bind these cores and an additional source of external confinement (e.g., magnetic field pressure) would be required. However, if one attempts to scale down the $^{13}$CO linewidths, according to the observations of high-density tracers in a small sample of sources, then most proto-stellar cores would result effectively gravitationally bound. ",
        "introduction": "Pre-stellar cores represent a very early stage of the star formation (SF) process, before collapse results in the formation of a central protostar. The physical and dynamical properties of these cores can reveal important clues about their nature. More than a thousand new pre- and proto-stellar cores have been identified by BLAST, the Balloon-borne Large-Aperture Submillimeter Telescope \\citep{pascale2008}, during its second long duration balloon (LDB) science flight in 2006. BLAST detected these cold cores in a $\\sim 50\\,$deg$^2$ map of the Vela Molecular Ridge (VMR) \\citep{netterfield2009}, and was capable of constraining the temperatures of the detected sources using its three-band photometry (250, 350, and 500\\,\\micron) near the peak of the cold core SED.  The VMR (\\citealp{mur:may}, \\citealp{Lis92}) is a giant molecular cloud complex within the galactic plane, in the area $260 \\degr \\la l \\la 272 \\degr$ and $-2 \\degr \\la b \\la 3 \\degr$, hence located outside the solar circle. Its main properties have been recently revisited by \\citet{netterfield2009} and \\citet{olmi2009}. In particular, \\citet{olmi2009} used both BLAST and archival data to determine the spectral energy densities (SEDs) and the physical parameters of each source detected by BLAST in the smaller region of Vela-D, where observations from IR (\\citealp{gianni07}) to millimeter wavelengths (\\citealp{massi07}) were already available.  In this work we will use the spectral line data of \\citet{elia2007} to perform an analysis of the dynamical state of the cores. In Section~\\ref{sec:dyn} we discuss which cores are gravitationally bound. Then in Section~\\ref{sec:pressure} we discuss the effects of pressure. We draw our conclusions in Section~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  In this paper we have analyzed the dynamical state of the population of the dense and cold cores detected by BLAST in the Vela-D molecular cloud, utilizing the sensitive maps at 250, 350 and 500\\,\\micron\\ obtained by BLAST during its 2006 LDB flight from Antarctica. We can summarize the results of this paper as follows, although some of our conclusions are still dependent on an exact measurement of the internal velocity dispersion of the cores:  1. The median $M_{\\rm core}/M_{\\rm BE} = 0.9$ for all sources in Vela-D, i.e. very near the critical value, assuming thermal support only. If we restrict ourselves to mass ranges for which we are progressively more complete, our sources exceed their Bonnor-Ebert mass by a median factor of $M_{\\rm core}/M_{\\rm BE} = 1.5 - 3.3$.  2. If the BLAST cores were purely isothermal structures, then almost all would be either confined by the external gas pressure, or by their gravity. According to the virial analysis of the cores with an unambiguous line-of-sight $^{13}$CO(2-1) spectrum, almost none of them would appear to be gravitationally bound and would also be dominated by non-thermal pressure. However, we think this is likely a consequence of the $^{13}$CO(2-1) linewidths overestimating the core internal velocity dispersion. Despite the uncertainties involved in this procedure, we have attempted to correct for this effect and found that $\\sim 30-50$\\% of the cores would then turn out to be gravitationally bound.   3.  The average total internal pressure of the cores appear to be higher than the ambient external pressure. If most cores were indeed gravitationally {\\it unbound} then self-gravity and surface pressure would be insufficient to bind the cores and form coherent structures. After application of the scale factor to the $^{13}$CO(2-1) linewidths, a significant fraction of the cores (mostly proto-stellar) would actually become gravitationally bound, and thus they would not need to be pressure-confined.  4. The cores that result gravitationally unbound even after application of the scale factor to the $^{13}$CO(2-1) linewidths, would still need an additional source of external confinement. If the confinement mechanism were an external magnetic pressure, the required $B$ would vary from a few tens to several hundreds of $\\mu$G. Measurements of the BLAST cores velocity dispersion is necessary to allow a more accurate dynamical analysis.  5. Alternatively, the gravitationally unbound cores could be transient structures, continuously produced in the region by turbulence."
    },
    "1005/1005.3795_arXiv.txt": {
        "abstract": "s{ The Pierre Auger Observatory has already collected more ultra high energy cosmic ray data than all the previous experiments. With an hybrid detection technique, it can provide coherent results on the flux, energy spectrum and arrival directions of the highest energy cosmic rays, and characterize the extensive air showers in order to probe the primary particle characteristics and its interactions. These results will be presented from the point of view of particle physics.} ",
        "introduction": "High energy cosmic rays may give important information on different areas of Physics. The sources where they are produced, the corresponding acceleration mechanisms, and even their precise composition in terms of primary nuclei remain unknown and are important questions from the astrophysical point of view. When they arrive at Earth, they collide with the atmosphere nuclei at center-of-mass energies which are orders of magnitude above the ones available in man-made particle accelerators, allowing the probe of particle physics at a new energy scale. In between, they propagate through the interstellar and intergalactic media, subjected to magnetic fields and to interactions with cosmic matter and radiation, having the possibility to give indirect indication about phenomena which arise only at large distances and energies.  The Earth atmosphere acts as an efficient calorimeter for the high energy cosmic rays. The initial particle interacts with the atmospheric nuclei creating large numbers of lower energy particles in a showering process. Nuclei will produce charged pions which can interact or decay, producing both the atmospheric neutrino flux and high energy muons that can be efficiently sampled at ground. Neutral pions will decay immediately into gammas feeding an electromagnetic shower which will carry most of the initial energy. The electromagnetic shower will produce isotropic fluorescence light by exciting the nitrogen molecules in the air which allows the imaging of the showers passing in front of a telescope regardless of their direction. The intensity of this light is directly proportional to the deposited energy. ",
        "conclusions": "The Pierre Auger Observatory has been working since the start of installation in 2004, and has now collected an amount of data corresponding to just two years of the full Observatory but larger than the predecessor detectors. Using a Hybrid technique, the Observatory profits from the direct energy determination from the Fluorescence Detectors, and extends that precision to the full Surface Detector statistics with virtually no dependence on simulations.  The fast decreasing high energy cosmic ray flux shows two features, which allow us to start attempting a coherent characterization of three regions: \\begin{itemize} \\item{below the {\\it ankle}, the flux dominated by protons created in our galaxy and trapped by its magnetic field;} \\item{above the {\\it ankle}, the flux is dominated by extra-galactic particles. These are not necessarily protons, on the contrary, their interactions in air seem to have higher cross-section and/or multiplicity. Neutral primaries have not been observed yet, and charged ones are deviated by the galactic magnetic fields leading to a featureless sky.} \\item{above $10^{19.5}$~eV, there is a rapid flux decrease. The energy is enough for the CMB to induce photo-disintegration of heavy nuclei and proton resonances, limiting the observable distance, but also enough for charged particles to overcame magnetic fields, allowing for charged particle astronomy in this limited part of the Universe.} \\end{itemize} However, more data and more analyzes are necessary to fully fix this simplified picture. While the identification of individual cosmic ray sources will mean an important breakthrough for Astrophysics, Particle Physics will be necessary to interpret what exactly arrives to Earth. Even if the data seem to indicate we are in the presence of heavy nuclei primary, no model can fully explain the data today. A combination of variables will give more information in the near future. The detailed study of these particle interactions in the atmosphere might constrain the hadronic models evolution to energies far beyond the reach of man-made accelerators."
    },
    "1005/1005.2192_arXiv.txt": {
        "abstract": "{Nuclear and starburst activity are known to often occur concomitantly. {\\it Herschel}-SPIRE provides sampling of the far-infrared (FIR) spectral energy distributions (SEDs) of type 1 and type 2 AGN, allowing for the separation between the hot dust (torus) and cold dust (starburst) emission. We study large samples of spectroscopically confirmed type 1 and type 2 AGN lying within the {\\it Herschel} Multi-tiered Extragalactic Survey (HerMES) fields observed during the science demonstration phase, aiming to understand their FIR colour distributions and constrain their starburst contributions. We find that one third of the spectroscopically confirmed AGN in the HerMES fields have 5$\\sigma$ detections at 250\\,$\\mu$m, in agreement with previous (sub)mm AGN studies. Their combined {\\it Spitzer}-MIPS and {\\it Herschel}-SPIRE colours (specifically $S_{250}/S_{70}$ vs $S_{70}/S_{24}$) quite clearly separate them from the non-AGN, star forming galaxy population, as their 24 $\\mu$m flux is dominated by the hot torus emission. However, their SPIRE colours alone do not differ from those of non-AGN galaxies. SED fitting shows that all those AGN need a starburst component to fully account for their FIR emission. For objects at $z>2$ we find a correlation between the infrared luminosity attributed to the starburst component, $L_\\mathrm{SB}$, and the AGN accretion luminosity, $L_\\mathrm{acc}$, with $L_\\mathrm{SB} \\propto L_\\mathrm{acc}^{0.35}$. Type 2 AGN detected at 250\\,$\\mu$m show on average higher $L_\\mathrm{SB}$ than type 1 objects but their number is still too low to establish whether this trend indicates stronger star formation activity.} ",
        "introduction": "Active galactic nuclei (AGN) and starburst activity, both among the most energetic extragalactic phenomena, have been studied separately for decades. However, it is only in the past decade and a half, with the advent of the infrared (IR) observatories (ISO, {\\it Spitzer} and now {\\it Herschel}; \\cite{pilbratt10}) that it has become clear that the two phenomena are related and, more often than not, happen concomitantly (e.g. \\cite{schweitzer06}).  The {\\it Spitzer} IRAC and MIPS cameras sampled the spectral energy distribution (SED) of extragalactic sources exactly where the peak of the AGN dust emission is expected to occur under the paradigm of an axisymmetric dust distribution (often referred to as {\\it torus}) surrounding the central super-massive black hole. The IRS spectrograph provided the necessary details of the silicate feature in emission and in absorption and we now have a better understanding of the physics of hot dust around AGN (e.g.\\ \\cite{hao07}; \\cite{levenson07}). The peak of the cold dust emission, however, a tracer of star formation, was beyond the wavelength range explored by {\\it Spitzer} for high redshift sources and it is only now, with the advent of SPIRE (\\cite{griffin10}), that the full FIR SEDs of galaxies can be built, all the way to 500\\,$\\mu$m.  The work presented here intends to build on the previous experience gathered with {\\it Spitzer} and BLAST. % We study large samples of spectroscopically confirmed type 1 and type 2 AGN in the largest {\\it Herschel} Multi-tiered Extragalactic Survey (HerMES\\footnote{http://hermes.sussex.ac.uk}; Oliver et al.\\ in prep.) fields observed during the science demonstration phase (SDP). The SEDs and IR properties of many of the objects in the samples have been studied in the past (e.g.\\ \\cite{richards06}; Hatziminaoglou et al.\\ 2008, 2009) and their AGN properties are well constrained. The idea is to extend such study to the larger wavelengths now observed by SPIRE in an effort to also constrain the starburst component of these objects. ",
        "conclusions": "Until recently, the study of type 1 and type 2 objects focused mainly on their X-ray-to-MIR emission, with the two extreme $\\lambda$ ranges sampling the nuclear activity (X-rays) and hot dusty torus (MIR) emission, respectively, and the middle wavelengths (optical-to-near IR) defining whether an AGN would be classified as type 1 (unobscured) or type 2 (obscured). With {\\it Herschel}-SPIRE we can now also accurately probe the cold dust component, heated by star formation, and look for differences between the types of active galaxies not only in the immediate environments of their nuclei but also in their hosts. The somewhat surprising lack of a strong correlation between the AGN intrinsic properties ($L_\\mathrm{acc}$) and the star formation activity as traced by $L_\\mathrm{SB}$, at least up to $z\\sim2$, possibly an effect of the narrow span of our sample in $L_\\mathrm{acc}$, requires further investigation with larger AGN samples that also cover a wider accretion luminosity range, including stacking sources undetected by SPIRE. The same holds for the seemingly higher star formation activity in type 2 AGN. Even though this may not necessarily conflict with geometric unification schemes, it certainly opens the way for investigations in new directions."
    },
    "1005/1005.5251_arXiv.txt": {
        "abstract": "We present AKARI near- to far-infrared images of the nearby edge-on spiral galaxy NGC~3079 in 10 photometric bands. The spectral energy distribution consists of continuum emission from dust with a single temperature of 28--33 K together with strong mid-infrared emission features from polycyclic aromatic hydrocarbons (PAHs). We derive the dust masses of $5.6 \\times 10^6 \\MO$ and $1.4 \\times 10^7 \\MO$ for the central 4 kpc region and the whole galaxy, respectively, and find that a gas-to-dust mass ratio is unusually high in the central region ($\\sim$ 1100) and even for the whole galaxy ($\\sim$ 860). The ratio of the surface brightness distribution at the wavelength of 7 $\\mathrm{\\mu m}$ to that at 11 $\\mathrm{\\mu m}$ suggests that the properties of PAHs have spatial variations. Emission from ionized and neutral PAHs is relatively strong in the center and the disk regions, respectively, suggesting stronger radiation field and thus relatively active star formation in the center. Yet the total infrared luminosities of the galaxy indicate rather low star formation rates. These results suggest that NGC~3079 is in an early-phase starburst stage. ",
        "introduction": "NGC~3079 is a highly-inclined ($i=84^\\circ$) edge-on spiral galaxy at a distance of 15.6 Mpc, the optical disk of which has a position angle of $166^\\circ$ (Irwin \\& Seaquist 1991; Sofue \\& Irwin 1992). The nucleus of NGC~3079 is classified as a low-ionization nuclear emission-line region (LINER; Heckman 1980) or Seyfert 2 (Ford et al. 1986). The galaxy is known to have an ultra-high-density molecular core ($3 \\times 10^8 \\MO$) in the center (Sofue et al. 2001). In addition, molecular gas with a total mass of $5 \\times 10^9 \\MO$ is contained within the central 2 kpc radius (Koda et al. 2002). The nucleus shows very bright $\\mathrm{H_2 O}$ emission, which is one of the most luminous $\\mathrm{H_2 O}$ maser sources (Henkel et al. 1984). Near the central region, kilo-parsec-scale outflows along the minor axis are observed at various wavelengths (radio continuum: Duric \\& Seaquist 1988; H$\\alpha$: Ford et al. 1986; Cecil et al. 2001; X-ray: Fabbiano et al. 1992). The superbubbles may have been produced by either a central active galactic nucleus (AGN, i.e. acretion of dense gas into the core; Irwin \\& Seaquist 1988; Hawarden et al. 1995) or a nuclear starburst (Sofue \\& Vogler 2001), but the exact origin is yet to be understood.  Infrared (IR) imaging observations of NGC~3079 are relatively few, in contrast to other wavelengths, mainly due to low spatial resolution especially in the far-IR. IR observations however provide very important tools to study dust components that trace star formation activities effectively. Young et al. (1989) and Klaas \\& Walker (2002) discussed the physical states of far-IR dust in NGC~3079 from the observations of IRAS and ISO, respectively. However the angular resolutions of IRAS and ISO are so large as compared to the apparent size of NGC~3079 ($\\timeform{7'.9} \\times \\timeform{1'.4}$; de Vaucouleurs et al. 1991) and therefore IRAS and ISO could not observe the galaxy as a spatially resolved far-IR source. Klaas \\& Walker (2002) derived the dust-to-gas mass ratio of 200 for the whole galaxy of NGC~3079, which is close to the accepted value of 100--200 for our Galaxy (Sodroski et al. 1997). Braine et al. (1997) and Stevens \\& Gear (2000) discussed the properties of dust in NGC~3079 on the basis of the spatially resolved submillimeter dust emission with IRAM and SCUBA, respectively.  Spitzer observed NGC~3079 with the IRAC, IRS, and MIPS. From the archived data, however, we find that the MIPS 160 $\\mathrm{\\mu m}$ and the IRAC 8 $\\mathrm{\\mu m}$ band images are heavily affected by saturation effects in the central region of NGC~3079. Weedman et al. (2005) and Bernard-Salas et al. (2009) presented mid-IR spectra of the nuclear region between 5 and 37 $\\mathrm{\\mu m}$ obtained by the IRS. These spectra showed that emissions from polycyclic aromatic hydrocarbons (PAHs) are dominant in the mid-IR. The strong PAH emission suggests that the central activity of NGC~3079 is better characterized by a nuclear starburst rather than an AGN.  In this paper, we present multi-band images of NGC~3079 in the near- to far-IR obtained by the AKARI satellite (Murakami et al. 2007). AKARI has two instruments, the Infrared Camera (IRC; Onaka et al. 2007) and the Far-Infrared Surveyor (FIS; Kawada et al. 2007), with which we obtain 10 photometric band images at central wavelengths of 3, 4, 7, 11, 15, 24, 65, 90, 140, and 160 $\\mathrm{\\mu m}$ by pointed observations. The 7 and 11 $\\mathrm{\\mu m}$ bands enable us to discuss the spatial distribution of the PAH emission. The 4 far-IR bands are efficient to determine the properties of interstellar dust. We perform spatially-resolved studies of the interstellar dust in the whole galaxy of NGC~3079. Here we present non-saturated images as well as new maps at 11 $\\mathrm{\\mu m}$ and 15 $\\mathrm{\\mu m}$, although the spatial resolutions are slightly worse than Spitzer due to the smaller diameter of the AKARI telescope (68.5 cm; Kaneda et al. 2005; 2007). ",
        "conclusions": "With AKARI, we obtain the near- to far-IR images of NGC~3079 in 10 photometric bands. We derive SEDs for the 4 kpc center region, the two 4 kpc disk regions, and the whole galaxy. The SEDs consist of far-IR continuum emission from dust with a single temperature of 28--33 K and strong mid-IR emission features from PAHs. We derive the dust masses of $5.6 \\times 10^6 \\MO$ and $1.4 \\times 10^7 \\MO$ for the central 4 kpc region and the whole galaxy, respectively, and find that NGC~3079 has unusually high dust-to-gas mass ratios in the central 4 kpc region ($\\sim$ 1100), and even for the whole galaxy ($\\sim$ 860). The ratio of the surface brightness distribution at the wavelength of 7 $\\mathrm{\\mu m}$ to that at 11 $\\mathrm{\\mu m}$ suggests spatial variations in the properties of PAHs; emission from ionized and neutral PAHs is relatively strong in the center and the disk regions, respectively. We detect extended emission from hot dust in the center, possibly associated with the superbubbles. Thus radiation field is strong and star formation is relatively active in the center. Yet the total infrared luminosities of the galaxy indicate rather low star formation rates of 2.6 $\\MO~\\mathrm{yr^{-1}}$ and 5.6 $\\MO~\\mathrm{yr^{-1}}$ for the central 4 kpc region and the whole galaxy, respectively. Hence we conclude that NGC~3079 is likely to be in an early phase of the starburst with a copious amount of dust-poor gas in the center region.  \\bigskip  We would like to thank all the members of the AKARI project for their intensive efforts. We also express many thanks to the anonymous referee for the useful comments. This work is based on observations with AKARI, a JAXA project with the participation of ESA."
    },
    "1005/1005.2537_arXiv.txt": {
        "abstract": "{Previous Spitzer and IRAS observations of the LMC suggest an excess of FIR emission with respect to the gas surface density traced by \\COT rotational emission lines and \\his 21 cm emission. This so-called ``FIR excess'' is especially noticeable near molecular clouds in the LMC, and has usually been interpreted as indicating the presence of a self-shielded H$_2$ component not traced by CO in the envelopes of molecular clouds. } {Based on Herschel HERITAGE observations taken as part of the Science Demonstration Phase, we examine the correlation between gas and dust surface densities at higher resolution than previously achieved.  We consider three additional possible causes for the FIR excess: X factor, FIR dust emissivity, and gas-to-dust ratio variations between the diffuse and dense phases of the ISM. } {We examine the structure of NT80 and NT71, two molecular clouds detected in the NANTEN \\COT survey of the LMC. Dust surface density maps were derived from the HERITAGE data. The gas phase is traced by MAGMA \\COT and ATCA$+$Parkes \\his 21 cm observations of the LMC. These data provide unprecedented resolution (1') to examine the structure of molecular clouds. The dust emissivity, gas-to-dust ratio, and X factor required to match the dust and gas surface densities are derived, and their correlations with the dust surface density are examined.  } {We show that the dust surface density is spatially correlated with the atomic and molecular gas phases. The dust temperature is consistently lower in the dense phase of the ISM than in the diffuse phase. We confirm variations in the ratio of FIR emission to gas surface density derived from \\his and CO observations. There is an excess of FIR emission, spatially correlated with regions of intermediate \\his and dust surface densities (A$_V$ = 1-2), and little or no CO. While there is no significant trend in the dust emissivity or gas-to-dust ratio with dust surface density, the X factor is enhanced at A$_V$ = 1-2. We conclude that H$_2$ envelopes not traced by CO and X factor variations close to the CO boundary may be more likely to cause these deviations between FIR emission and gas surface density than gas-to-dust ratio or emissivity variations.} {} ",
        "introduction": "\\indent Dust, neutral atomic hydrogen (\\hi), and molecular hydrogen (H$_2$) are the prime constituents of the interstellar medium in galaxies out of which stars form, but their amounts are often poorly known. In dense clouds, dust shields both H$_2$ and its tracer CO from dissociation by the ambient interstellar radiation field (ISRF). Unlike CO, H$_2$ is also strongly self-shielding. In the solar neighborhood,  H$_2$ forms at $A_V$ $\\geq$ 0.14, while CO requires $A_V$ $\\geq$ 0.8 \\citep{wolfire10}.  Molecular clouds (MCs) thus consist of dense cores where CO and H$_2$ coexist and less dense envelopes of H$_2$ with little or no CO. In lower-metallicity environments with strong irradiation, the poorly shielded CO fills a much smaller fraction of the H$_2$ volume. In those galaxies use of a standard conversion factor X$_{\\mathrm{CO}}$ to estimate H$_2$ column densities from observed CO emission causes large amounts of H$_2$ to be missed \\citep[see e.g.,][]{glover10}.\\\\ \\indent The nearest low-metallicity galaxies are the Large Magellanic Cloud (LMC) and the Small Magellanic Cloud (SMC) with [C] and [O] abundances 0.25/0.50 and 0.10/0.25 relative to solar abundances \\citep{pagel03} and distances of 50 kpc \\citep{schaefer08} and 62 kpc \\citep{szewczyk09} respectively. The FIR emission from dust has been used to establish that indeed much H$_2$ is not traced by CO and exhibits a so-called ``FIR excess'', implying X factors 3-6 (LMC) and 20-60 (SMC) times higher than in the solar neighborhood \\citep{israel97, leroy07, leroy09, bernard08}. Their analysis assumes that dust grain emissivity and gas-to-dust ratios are the same in dense H$_2$ clouds and more tenuous \\his regions. Our goal in this Letter is to explore whether these assumptions are justified. To this end, we examine the structure of two MCs in the LMC, NT80 and NT71 \\citep{fukui08}. Both clouds are relatively quiescent \\citep{kawamura09}, with star formation rates implied by H$\\alpha$ and 24 $\\mu$m emission of 0.018 and 0.042 M$_{\\odot}$/kpc$^2$/yr. NT80 is located in a direction practically devoid of H$\\alpha$ emission, while NT71 is associated with the faint filamentary H$\\alpha$ nebula DEM 110 \\citep{davies76}. We examine the correlation between dust and gas based on {\\it HERschel Inventory of The Agents of Galactic Evolution} (HERITAGE) data and MAGellanic Mopra Assessment (MAGMA, PI. T. Wong) \\COT data. These data provide unprecendented resolution (15 pc) to observe the structure of MCs in the LMC. ",
        "conclusions": "\\indent We have examined the correlation between dust, atomic, and molecular gas using HERITAGE, ATCA \\his 21 cm, and MAGMA CO observations of two LMC molecular clouds. The dust temperature appears consistently lower in the dense phase than in diffuse regions. The dust surface density is spatially correlated with the atomic and molecular phases, making {\\it Herschel}'s angular resolution and complete coverage of the IR SED a powerful tracer of molecular gas. We have however observed an excess of FIR emission with respect to the gas surface density implied by CO and \\his observations, which occurs at intermediate dust surface densities (0.1-0.2 M$_{\\odot}$/pc$^2$), outside and close to the CO boundary. This likely indicates that molecular clouds are surrounded by envelopes of H$_2$ not traced by CO. The presence of unaccounted for H$_2$ envelopes is further supported by an increase in the X factor at intermediate dust surface densities, corresponding to the transition region between dissociated and shielded CO. \\\\ \\indent We reviewed two alternative explanations for the FIR excess: variations in dust emissivity and the gas-to-dust ratio between the diffuse and dense phases of the ISM. We derived the dust emissivity and gas-to-dust ratio required to match the observations, and examined their correlations with the dust surface density in order to evaluate the plausibility of each hypothesis. We found that the dust emissivity and gas-to-dust ratio in NT71 and NT80 are constant with $\\Sigma_{\\mathrm{dust}}$ within the scatter, and conclude that dust emissivity and gas-to-dust ratio variations are therefore unlikely to be responsible for the FIR excess observed near these clouds. Variations in emissivity and gas-to-dust ratio between the dense and diffuse ISM phases cannot be definitively ruled out however, due to uncertainties at low dust surface density that are caused by offsets in the zero levels of the \\his and dust maps. In the immediate future, we will conduct a full investigation of all these effects using detailed modeling in combination with the completed HERITAGE survey of both Magellanic Clouds."
    },
    "1005/1005.5317_arXiv.txt": {
        "abstract": "We studied the effect of the perturbation of the meridional flow in the activity belts detected by local helioseismology on the development and strength of the surface magnetic field at the polar caps. We carried out simulations of synthetic solar cycles with a flux transport model, which follows the cyclic evolution of the surface field determined by flux emergence and advective transport by near-surface flows. In each hemisphere, an axisymmetric band of latitudinal flows converging towards the central latitude of the activity belt was superposed onto the background poleward meridional flow. The overall effect of the flow perturbation is to reduce the latitude separation of the magnetic polarities of a bipolar magnetic region and thus diminish its contribution to the polar field. As a result, the polar field maximum reached around cycle activity minimum is weakened by the presence of the meridional flow perturbation. For a flow perturbation consistent with helioseismic observations, the polar field is reduced by about 18\\% compared to the case without inflows. If the amplitude of the flow perturbation depends on the cycle strength, its effect on the polar field provides a nonlinearity that could contribute to limiting the amplitude of a Babcock-Leighton type dynamo. ",
        "introduction": "Surface flux transport models treat the evolution of the large-scale magnetic field on the surface of the Sun \\citep[e.g.,][]{Wang:etal:1989b, Schrijver:2001, Mackay02, Baumann04}. In such models, the evolution of the radial magnetic field at the solar surface is governed by the emergence of new flux in the form of bipolar magnetic regions and by advective transport through large-scale flows (differential rotation, meridional circulation) and supergranular turbulent diffusion. The well-known cyclic variations of differential rotation in the form of zonal flows \\citep[e.g.,][]{Howard:Labonte:1980, Howe:etal:2006b} so far have not been considered in flux transport simulations.  On the other hand, variations in the large-scale meridional flow \\citep{Komm93, Basu03,Hathaway:Rightmire:2010} have been considered in flux-transport simulations by assuming cycle-to-cycle changes in the overall amplitude of the flow \\citep{Wang02_b, Dikpati04, Wang:etal:2009}.  Another cycle-related modulation of the surface flow field is the modulation of the axisymmetric component of the meridional flow in the form of bands of latitudinal velocity centered on the dominant latitudes of magnetic activity, first detected at depths greater than 20~Mm \\citep {Chou:Dai:2001, Beck:etal:2002}. In the case of near-surface flows, the residual meridional flow velocities (after subtraction of the mean flow) during cycle 23 were of the order of $\\pm $3$-$5~m$\\,$s$^{-1}$ and converge toward the dominant latitudes of magnetic activity while migrating towards the equator in parallel to the activity belts \\citep{Gizon08, Gonzalez08, Gonzalez:etal:2010}.  These flows are probably related to the meridional motions of sunspots and pores \\citep[][and references therein]{Ribes:Bonnefond:1990} and other magnetic features \\citep{Komm:1994, Meunier:1999}. The cumulative effect of the near-surface horizontal flows converging towards active regions \\citep[e.g.,][]{Haber:etal:2004, Hindman:etal:2004} appear to contribute to the axisymmetric meridional flow perturbation \\citep{Gizon:2004, Gonzalez08}, but there is evidence that at least part of this perturbation is unrelated to surface activity \\citep{Gonzalez:etal:2010}.  The effect of the near-surface inflows on the evolution of single active regions was recently studied by \\citet{De_Rosa06}. Considering results obtained with a surface flux transport model \\citep[cf.][]{Schrijver:2001}, these authors find that inflows of the order of $\\sim10$~m$\\,$s$^{-1}$ significantly affect the dispersal of magnetic flux from an isolated active region. These results indicate that the axisymmetric meridional flow perturbations associated with the activity belts could also affect the evolution of the solar surface field on a global scale. Particularly interesting in this connection is the effect on the polar field strength, which is an important source of the heliospheric field and also plays a significant role in Babcock-Leighton-type dynamo models. Here we present results of solar-cycle simulations using the flux transport code of \\citet{Baumann04}, including axisymmetric bands of converging latitudinal flows centered on the migrating activity belts.   The aim of this work is to study the general effect of these flows on the evolution of the solar surface field, and particularly on the strength of the polar field. This is an exploratory study focussing on understanding the physical mechanisms; we do not intend to reproduce any actual solar data. We need not consider the zonal flows in this study because the buildup of magnetic field to the Sun's poles is dominated by the latitude separation of the polarities of a bipolar magnetic region and thus essentially is an axisymmetric problem \\citep{Cameron07}; zonal flows \\citep[and differential rotation in general, see][]{Leighton:1964} have no effect on the amount of signed flux reaching the poles.  This paper is organized as follows. The flux transport model is described in Section~\\ref{sec:model}. The relevant effects of the latitudinal flow bands on the surface flux evolution are illustrated with simulations of single bipolar regions in Section~\\ref{sec:bipoles}. The results of full solar-cycle simulations are presented in Section~\\ref{sec:cycles}, which includes a study of the dependence of the polar field on various model parameters. The implication of our results are discussed in Section~\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  The results presented here show that the observed cycle-related meridional flow perturbations in the form of bands migrating with the activity belts decrease the strength of the polar fields resulting from the latitudinal transport of surface flux. For a flow perturbation corresponding to the helioseismic observations, this reduction amounts to about 18\\% with respect to the case without flow perturbation. This indicates that these effects should be taken into account in surface flux transport simulations aiming at a quantitative pre- or postdiction of the polar field strength.  It is doubtful whether this kind of flow perturbation could have significantly contributed to the low polar polar field strength during the activity minimum between solar cycles 23 and 24 \\citep[e.g.,][]{Schrijver:Liu:2008} as compared to previous minima: the perturbation is probably present during every cycle, so that only an {\\em increase} of the perturbation in cycle 23 compared to its amplitude in previous cycles would contribute to a comparatively weaker polar field during the recent minimum. In any case, other effects must have been affecting the polar field in addition since the observed reduction by nearly a factor of 2 exceeds the decrease that could be caused by the flow perturbation considered here.  The observed variation of the flow perturbation amplitude during the activity cycle \\citep{Gonzalez:etal:2010} and the probable contribution of the near-surface inflows toward active regions to the driving of the perturbations \\citep {Gizon08} suggest that the amplitude of the flow perturbation should increase with cycle strength. According to our results, this would lead to a stronger reduction of the polar field built up during cycles of higher activity. Since, in the framework of a Babcock-Leighton dynamo, the polar field is a measure of the poloidal field providing the basis for the toroidal field in the subsequent cycle, the meridional flow perturbation is potentially important for the nonlinear modulation and limitation of the cycle amplitude. Furthermore, we have also seen that the polar field decreases for increasing starting latitude, $\\lambda_0$, of the activity belt at the beginning of a cycle. Since stronger cycles typically have higher values of $\\lambda_0$ \\citep{Solanki:etal:2008}, this relationship would strengthen the nonlinear effect of the flow perturbation in the subsequent cycle amplitude.  We conclude that, in addition to global variations of the meridional flow speed \\citep{Wang02_b, Wang02}, the cyclic perturbation of the meridional flow by converging bands migrating with activity belts has an appreciable effect on the build-up of the magnetic field at the polar caps. Its relation to the strength of a cycle means that the flow perturbation could be an important factor in determining the amplitude of Babcock-Leighton-type flux transport dynamos."
    },
    "1005/1005.1702_arXiv.txt": {
        "abstract": "The occurrence of digits one through nine as the leftmost nonzero digit of numbers from real world sources is often not uniformly distributed, but instead, is distributed according to a logarithmic law, known as Benford's law. Here, we investigate systematically the mantissa distributions of some pulsar quantities, and find that for most quantities their first digits conform to this law. However, the barycentric period shows significant deviation from the usual distribution, but satisfies a generalized Benford's law roughly. Therefore pulsars can serve as an ideal assemblage to study the first digit distributions of real world data, and the observations can be used to constrain theoretical models of pulsar behavior. ",
        "introduction": "Pulsars, celestial lighthouses in the sky, are excellent natural laboratories for the research of fundamental properties of matter under the circumstances of strong gravity, strong magnetic field, high density, and extremely relativistic condition. Pulsar physics has been a forefront field of both astronomy and physics for more than 40 years since its discovery. Though exciting progresses in this domain have significantly enlarged our knowledge of astronomical environment and physical processes, there still remain outstanding problems~\\cite{gs06araa,wnrs07ppnp}. Nowadays, due to the efforts made by modern observational instruments covering over various wavelength scopes, \\emph{e.g.}, ROSAT, BeppoSAX, Chandra, XMM-Newton, HST, and VLT, people accumulated a sea of pulsar data, and the data are accreting significantly all these days. Hence data analysis and statistical synthesis become a vital step to characterize pulsar properties and reveal their inner regularities~\\cite{l04,bkps07}.   In this paper, we perform a systematic investigation of the mantissa distributions of pulsars for the first time. The mantissa $m \\in (-1,-0.1~]\\cup [~0.1,~1)$ is the significant part of a floating-point number $x$, defined as $x = m \\times 10^n$, where $n$ is an integer. In this paper, if not noted explicitly, we always postulate that the numbers are positive for succinct statement.  One might presume that the mantissas of any randomly chosen data set are approximately uniformly distributed, but that is not the case in real world. Instead, as stated by Newcomb~\\cite{n81ajm}, \\emph{``the law of probability of the occurrence of numbers is such that all mantissae of their logarithms are equally likely''}. Subsequently, it leads to the conclusion that the first significant digit, i.e., $1,2,...,9$, of mantissa is logarithmically distributed, where the number $1$ appears almost seven times more often than that of the number $9$. The probability of the occurrence of the first digit can be expressed in an analytical formula, called Benford's law~\\cite{b38paps} after the name of its second discoverer, \\begin{equation}\\label{benford} P(k) = \\log_{10} (1 + \\frac{1}{k}), \\, k=1,2,...,9 \\end{equation} where $P(k)$ is the probability of a number having the first digit $k$.  Empirically, the areas of lakes, the lengths of rivers, arabic numbers on the front page of a newspaper~\\cite{b38paps}, physical constants and distributions~\\cite{bk91ajp,lm}, the stock market indices~\\cite{l96as}, file sizes in a personal computer~\\cite{tfgs07ejp}, survival distributions~\\cite{lse00as}, widths of hadrons~\\cite{sm09mpla}, even dynamical systems~\\cite{tbl00c,bbh05tams,b05dcds}, conform to the peculiar law well. Nevertheless, there also exist other types of data, e.g., lottery and telephone numbers, which do not obey the law. Unfortunately, there is no a priori criteria yet to judge which type a data set belongs to. In practice, the law is already applicable in distinguishing and ascertaining fraud in taxing and accounting~\\cite{n96jata,nm99a,rr03ja,gw04cssc}, and speeding up calculation and minimizing expected storage space in computer science~\\cite{bb85c,s88ipce,bhs07amm}.  Since its second discovery in 1938, many attempts have been tried to explain the underlying reason for Benford's law. For theoretical reviews, see papers written by Raimi~\\cite{r69amm,r76amm,r85paps} and Hill~\\cite{h95amm,h95pamc,h95ss,h98as,hs05spa}. Nowadays, many breakthrough points have been achieved in this domain, though, there still lacks a universally accepted final answer. In mathematics, Benford's law is the only digit law that is scale-invariant~\\cite{bhm08jtp}, which means that the law does not depend on any particular choice of units, discovered by Pinkham~\\cite{p61ams}. Also Benford's law is base-invariant~\\cite{h95amm,h95pamc,h95ss}, which means that it is independent of the base $d$ you use. In the binary system ($d$=2), octal system ($d$=8), or other base system, the data, as well as in the decimal system ($d$=10), all fit the general first digit law, $P(k) = \\log_{d}(1+1/k)$, $k \\in \\{1, 2, ..., d-1\\}$. Theoretically, Hill proved that \\emph{``scale-invariance implies base-invariance''}~\\cite{h95amm} and \\emph{``base-invariance implies Benford's law''}~\\cite{h95pamc} mathematically in the framework of probability theory. He also proved that random entries, picked from random distributions, form a sequence whose significant digit distribution approaches to Benford's law~\\cite{h95ss}.  Intending to uncover some regularities of pulsar data, and also to explore new domains of the digit law, we investigate the mantissa distributions of pulsar quantities systematically. We find that the first nonzero digit of mantissas displays unevenness according to the logarithmic law. The exceptions are the barycentric period and rotation frequency, which show significant deviations from Benford's law. Further, we also discuss various properties of the digit law, and perform the generalized Benford's law to data sets of barycentric period and rotation frequency of pulsars as well. Therefore the data of pulsars provide an ideal assemblage for further studies on the first digit law of the nature. ",
        "conclusions": "In this paper, we present systematic analysis on the first digit distributions of mantissas of most fundamental quantities of pulsars, including barycentric period and rotation frequency, together with their time derivatives, and many more. The results reveal obvious departures from the uniform distribution, and small digits are more prevalent than large ones according to a logarithmic formula, called Benford's law. However, not all data sets conform to it. Artificial and restricted data sets often diverge from the law, such as the expected minimum and median companion masses of pulsars in our study. Furthermore, % the generalized Benford's law is applied to the barycentric period and rotation frequency of pulsars.  Since the first discovery of the peculiar significant digit law more than one hundred years ago, many crucial understandings of the inner regularities from mathematical viewpoint have been achieved. But there still remain big challenges to scientists to look into more profound physical reasons of such remarkable property of the nature. Thus the revelation of the hitherto unnoticed regularities in the pulsar data opens a new window to look into novel aspects of astronomical objects."
    },
    "1005/1005.1534_arXiv.txt": {
        "abstract": "A detailed observation of microglitch phenomenon in relatively slow radio pulsars is presented. Our analyses for these small amplitude jumps in pulse rotation frequency ($\\nu$) and/or spin down rate ($\\dot{\\nu}$) combine the traditional manual detection method (which hinges on careful visual inspections of the residuals of pulse phase residuals) and a new, and perhaps more objective, automated search technique (which exploits the power of the computer, rather than the eyes, for resolving discrete events in pulsar spin parameters). The results of the analyses of a sample of 26 radio pulsars reveal that: (i) only 20 pulsars exhibit significant fluctuations in their arrival times to be considered suitable for meaningful microglitch analyses; (ii) a phenomenal 299 microglitch events were identified in $\\nu$ and/or $\\dot{\\nu}$: 266 of these events were found to be simultaneously significant in $\\nu$ and $\\dot{\\nu}$, while 19 and 14 were noticeable only in $\\nu$ and $\\dot{\\nu}$, respectively; (iii) irrespective of sign, the microglitches have fractional sizes which cover about 3 orders of magnitude in $\\nu$ and $\\dot{\\nu}$ ($10^{-11} < |\\Delta{\\nu}/\\nu| < 2.0\\times10^{-8}$ and $5.0\\times10^{-5} < |\\Delta{\\dot{\\nu}}/\\dot{\\nu}| < 2.0\\times10^{-2}$) with median values as $0.78\\times10^{-9}$ and $0.36\\times10^{-3}$, respectively. ",
        "introduction": "Discrete discontinuities in the rotation of pulsars are broadly of two sorts: macroglitches and microglitches (e.g.\\ Cordes, Downs \\& Krause-Polstorff 1988). Macroglitches (conventionally known as glitches) are characterised by a sudden increase in pulsar rotation frequency, $\\nu$, usually accompanied by an increase in magnitude of the spin-down rate, $\\dot{\\nu}$, (e.g.\\ Urama 2002; Wang et al.\\ 2001; Lyne et al.\\ 1996; Flanagan 1990; Lyne 1987; Lyne \\& Pritchard 1987). Macroglitches are arguably the more spectacular events, characterised by fractional changes in pulsar spin frequency ($\\Delta{\\nu}/\\nu$) and spin-down rate ($\\Delta{\\dot{\\nu}}/\\dot{\\nu}$) in the range of $\\sim$ 10$^{-8}$ $-$ 10$^{-6}$ and 10$^{-4}$ $-$ 10$^{-1}$, respectively, (Lyne, Shemar \\& Smith 2000; Shemar \\& Lyne 1996). To date, about 280 macroglitches have been reported in about 100, mostly young, pulsars (Melatos, Peralta \\& Wyithe 2008 and references therein). Macroglitches have been extensively studied and are now firmly associated with some unique features. The jumps in the rotation frequency ($\\Delta{\\nu}$) and the spin-down rate ($\\Delta{\\dot{\\nu}}$) have a definite signature, ($\\Delta{\\nu}$, $\\Delta{\\dot{\\nu}}$) = (+, $-$), corresponding to increases in the pulsar spin rates and magnitude of the spin-down rates (Melatos et al.\\ 2008 and references therein). It is worthy of note that positive $\\Delta{\\dot{\\nu}}$ have been reported in a few macroglitches (with $\\Delta{\\nu}/\\nu > 10^{-7}$: see, for example, Middleditch et al.\\ 2006; Urama \\& Okeke 1999 and references therein). However, we should also mention that the error estimates in the measured values of $\\Delta{\\dot{\\nu}}/\\dot{\\nu}$ in all these cases are, uncomfortably, too large. For most well observed macroglitches, the jumps in the rotation frequency are accompanied by periods of relaxation towards the unperturbed pre-glitch values, which vary on a wide range of timescales (Lyne et al.\\ 1996; Flanagan 1995, 1990; Cordes et al.\\ 1988; Lyne 1987; Lyne \\& Pritchard 1987). Macroglitches are now widely believed to originate from some form of dynamical change within neutron star interior, such as a sudden and irregular transfer of angular momentum from the more rapidly rotating inner superfluid components to the slowly spinning crust (Ruderman, Zhu \\& Chen 1998; Alpar et al.\\ 1996, 1984)  On the other hand, microglitches constitute a class of small amplitude, but resolvable jumps in both, or either of the pulsar rotation frequency and its first time derivative. Unlike the spectacular macroglitch events, the study of microglitches have received a far less attention in the last three decades of pulsar timing observations. As a result, very little is known about this class of rotational discontinuities (Chukwude 2002; Cordes et al.\\ 1988). To date, only about 90 microglitches have been reported in $\\sim$ 20 radio pulsars (Cordes \\& Downs 1985, hereafter CD85; Cordes et al.\\ 1988; D'Alessandro et al.\\ 1995, hereafter DA95). The amplitudes of the jumps vary over a wide range, but are generally believed to be within $\\Delta{\\nu}/\\nu < 10^{-10}$ and $\\Delta{\\dot{\\nu}}/\\dot{\\nu} < 10^{-3}$. However, unlike macroglitches, where the jumps are associated with a definite signature, microglitches have all possible combination of signs and are extraordinarily difficult to understand in the framework of all known macroglitch theories (D'Alessandro 1997; Cordes et al.\\ 1988; CD85).  Previous observations of microglitches in radio pulsars have relied solely on visual inspection of carefully calculated residuals of pulse rotation frequency and its first time derivative ($\\Delta{\\nu}$ and $\\Delta{\\dot{\\nu}}$, respectively). However, observational constraints have permitted only indirect calculations of $\\Delta{\\nu}$ and $\\Delta{\\dot{\\nu}}$ via successive numerical differentiation of phase residuals obtained from 2nd-order model fits to barycentric times of pulse arrival (DA95 Cordes et al.\\ 1988; CD85). In principle, scatters in the differentiated phase residuals that are significantly in excess of measurement uncertainties are attributed to discrete rotational discontinuities. The effectiveness of this technique for microglitch observation has, hitherto, been plagued with sparsely sampled timing data. This results in a paucity of data, which introduces significant ambiguities in both the epoch and amplitude of candidate events. For instance, the process of numerically differentiating pulse phase residuals could require data segments with widely varying lengths. Hence, candidate microglitch events whose epochs are located deep within the segments could easily be averaged out during the process, thereby reducing the overall sensitivity of the technique for microglitch observation.  A more comprehensive observational description of the phenomenon of microglitches has, however, become an indispensible part of the current quest for improved statistics of radio pulsar microglitches. This is particularly important for the much sought-after better insight into the incidence and other features of these small amplitude discontinuities in pulsar rotation. Specifically, such a description is required in order to ascertain the applicability of the various models involving the neutron star interior and/or magnetospheric torque fluctuations (e.g.\\ Glampedakis \\& Andersson 2009; Melatos \\& Warszawski 2009; Warszawski \\& Melatos 2008; Melatos \\& Peralta 2007; Peralta et al.\\ 2006; Link, Epstein \\& Baym 1993; Alpar et al.\\ 1984; Cordes \\& Greenstein 1981; Arons 1981). Some potential benefits of this include a better understanding of the possible origin of microglitches and, perhaps, the relationship, if any, between macroglitches and microglitches.  In this paper, we present an in-depth observation of the phenomenon of microglitches in a sample of 26 radio pulsars. Our data, which span about 16 years for each object, are characterised by much improved sampling intervals of $\\sim$ 0.05 to  40 d. The short observation intervals and the resultant high data density allowed us to introduce alternative and, perhaps, more objective methods of resolving small amplitude discrete rotational discontinuities in radio pulsars. {\\bf In a forthcoming paper, we will compare the microglitching rates and size distributions with the glitch results recently published by Melatos et al.\\ (2008).}  \\section[]{OBSERVATIONS} Regular timing observations of all objects in the current sample of  26 radio pulsars commenced at Hartebeesthoek Radio Astronomy Observatory between 1984 January and 1987 May and are still ongoing. However, only the timing data accumulated up to 2002 October are reported in this paper. HartRAO regular pulsar timing observation was, however, interrupted between 1999 June and 2000 August due to a major hardware upgrade. Only the pulsars B0833$-$45 and B1641$-$45, on real time glitch monitoring program, were scheduled during this period. The two pulsars are, nontheless, excluded from current analysis. In HartRAO, pulse times of arrival (TOAs) were measured regularly at intervals $\\sim$ 1 $-$ 14 days using the observatory 26-m parabolic radio telescope. Pulses from radio pulsars were recorded by a single 10 MHz bandwidth receiver centred near 13  or 18 cm and no pre-detection dedispersion hardware was implemented during the period under consideration.  For each pulsar, detected pulses are smoothed with an appropriate filter-time constant, and folded over $N_{\\mathrm{p}}$ consecutive rotation periods to beat down the background noise.  $N_{\\mathrm{p}}$ is different for different pulsars, but it generally varies between 500 and 5000 for the current sample. This corresponds to an integration time in the range of $\\sim$ 48 s $-$ 32 mins. An integration was usually started at a particular second by synchronization to the station clock, which is derived from a hydrogen maser and is referenced to the Universal Coordinated Time (UTC) via the Global Position satellite (GPS). On average, three such on-line integrations were made for each pulsar during a typical observing session. Details of data acquisition and reduction at HartRAO have been described elsewhere (Flanagan 1995).  \\section[]{TIMING ANALYSIS AND RESULTS} The resulting topocentric arrival times were transformed to infinite observing frequency at the Solar System Barycentre (SSB) using the Jet Propulsion Laboratory DE200 solar system ephemeris (Standish 1982) and TEMPO2 software package (Hobbs, Edwards \\& Manchester 2006). Subsequent  modelling of the barycentric times of pulse arrival (hereafter referred to as BTOAs) and further analyses were accomplished with the HartRAO in-house pulsar timing analysis software (CPHAS), which is based on the standard pulsar timing technique (e.g.\\ Manchester \\& Taylor 1977) and is well described in Flanagan (1995). To account for the pulsar deterministic spindown, the BTOAs were modelled with a simple Taylor series of the form (e.g.\\ Lorimer \\& Kramer 2005) \\begin{equation} \\phi(t) = \\phi_{0} + \\nu(t-t_{0}) + \\frac{1}{2}\\dot{\\nu}(t-t_{0})^{2}, \\end{equation} where $\\phi_{\\mathrm{0}}$ is the phase at an arbitrary time $t_{\\mathrm{0}}$, $\\nu$ and $\\dot{\\nu}$ are, respectively, the rotation frequency and its first time derivative. Equation~(1) presumes that pulsar deterministic spindown follows a simple power-law relation of the form $\\dot{\\nu} = K\\nu^{n}$: where $K$ is a positive constant which depends on neutron star magnetic moment and moment of inertia and $n$ is the torque braking index. For the simplest standard model, in which the torque braking processes are largely due to vacuum magnetodipole radiation at the pulsar rotation frequency (e.g.\\ Goldreich \\& Julian 1969; Pacini 1968) $n = 3$. The difference between the observed BTOAs and the predictions of the best-fit model for pulsar rotation (the phase or, more appropriately, the time residuals) are shown in Fig.~1 for a selection of six HartRAO pulsars. The observed phase residuals of six pulsars (B0450$-$18, B1133+16, B1426$-$66, B1451$-$68,  B1933+16 and B2045$-$16) are dominated by intrinsic scatters of unusual large amplitudes. {\\bf The intrinsic scatter refers to a combination of flux noise and some other additional noise: which are largely due to effects such as pulse-to-pulse phase jitter, which is intrinsic to the pulsar, and interstellar scintillation, which causes fluctuations in the observed pulse strength within the integration time and bandwith (Cordes \\& Downs 1985). Scatters in BTOA residuals dominated by these noise processes more or less have white noise statistics, irrespective of their amplitudes (e.g.\\ Chukwude 2002).} The sizes of the BTOAs scatters suggest that they could have completely swallowed up any genuine, smaller amplitude timing noise activity inherent in these objects. For these peculiar pulsars, the rms timing noise ($\\sigma_{\\mathrm{TN}} = \\sqrt{\\sigma_{\\mathrm{R}}^{2} - \\sigma_{\\mathrm{W}}^{2}}$: where $\\sigma_{ \\mathrm{R}}$ is the root-mean-squares phase residuals from 2nd-order polynomial (i.e.\\ a model[$\\nu$, $\\dot{\\nu}$]) fit to the entire data span and $\\sigma_{\\mathrm{W}}$ is the rms white noise) is less than 5 mP.  Estimates of $\\sigma_{\\mathrm{W}}$ were obtained using pairs of BTOA residual data, from 2nd-order models, whose separations do not exceed 1 day (see, e.g.\\ Chukwude 2002). The signal-to-noise ratio (SNR = $\\sigma_{\\mathrm{R}} / \\sigma_{\\mathrm{R}}$) of these objects are typically less than 3 and they are excluded from further microglitch analysis in this paper. For the remainder of 20 pulsars, the observed SNR lie between $\\sim$ 8 and 700, suggesting a wide dispersion in the level of the observed timing activity. The phase residuals of some of the objects (e.g.\\ B0740$-$28, B1323$-$62, B1356$-$60 and B1749$-$28) are characterised by series of sudden slope changes. Such discontinuous rotational behaviour have been widely attributed to enhanced activity of small-amplitude discrete rotational jumps (Chukwude 2002; D'Alessandro \\& McCulloch 1997; CD85).  \\begin{figure} \\includegraphics[width=77mm]{p1fig1.eps} \\caption{The phase residuals, defined in the sense of model-predicted minus observed BTOAs, for a selection of 6 HartRAO radio pulsars over the interval between 1984 and 1999. The pulsar B0450$-$18 typifies objects whose BTOA residuals over the $\\sim$ 13-yr period are dominated by large amplitude scatter, which have predominantly white noise statistics. Meaningful microglitch analyses were, generally, not possible for this category of pulsars using current techniques. } \\end{figure}  \\subsection{BTOA Error Assignment} Assignment of correct errors to the pulse times of arrival data is crucial for a meaningful microglitch analysis. Such analysis usually involves a careful calculation of numerical values of  small amplitude jumps in pulsar spin parameters, the associated formal errors and the determination of the significance of the jump parameters relative to the formal errors. Therefore, an accurate BTOA error assignment improves both the sensitivity of the technique and the reliability of the results. In order words, under- or over-estimation of the errors assigned to BTOAs  would seriously affect the results of any glitch analysis, especially the microglitches. Correct assignment of BTOA errors is particularly important for the newly introduced automated microglitch detection technique. This technique employs the goodness of fit parameter, which relies heavily on correctly weighted least-square fits, as a tracer of microglitch events.  Unfortunately, the formal BTOA errors returned by the HartRAO in-house routines used to extract arrival times from observations are known to be grossly underestimated (Flanagan 1995). Consequently, realistic BTOA errors are obtained from the real scatter in the BTOA data (Flanagan 1995). Short segments of data, spanning between $\\sim$ 50 and 200 days, are modelled with 2nd-order polynomial plus a dispersion measure (DM) term. The length of the segments, which depended on both the timing activity level and the density of BTOAs, is such that the resulting phase resdiduals are statistically equal to zero. Hence, the segment lengths are expected to be shorter for pulsars with enhanced microglitch activity or near epochs of macroglitches to compensate for the expected additional scatter. The residuals are manually examined for possible outliers, which could bias the error estimates. Finally, BTOA errors are calculated from the real scatter in the phase residuals. Our estimator (following Flanagan 1995) is the two-sample variance calculated over successive blocks of length $\\Delta T$: where $\\Delta T \\sim$ 5 $-$ 10 days, depending on the data sampling rates. The error assigned to each BTOA is the rms of the two-sample variance of the block within which it falls. Although this method, unfortunately, results in BTOA data errors not being strictly independent, this problem is somewhat alleviated since our model fits involve data lengths $T$ (where $T \\gg \\Delta T$). The 13 and 18 cm observations were separately analysed, since the amplitudes of the observed scatter are different at the two frequencies. BTOA errors calculated using this method are found to be larger than those returned by the fitting routine, on average, by factors of $\\sim$ 2 and 4 for the 18 and 13 cm data, respectively. It has been shown (Chukwude 2007; Urama 2002; Flanagan 1995) that BTOA uncertainties obtained in this manner represent more realistic estimates.  \\begin{figure} \\includegraphics[width=75mm]{p1fig2.eps} \\caption{Time evolution of the frequency residuals ($\\Delta{\\nu}(t)$), calculated as described in the text, for a selection of 4 HartRAO pulsars over the period between 1984 and 2001. Arrows indicate the points where scatters in $\\Delta{\\nu}$ are signficantly ($>$ 2-$\\sigma$) larger than the measurement uncertainties, suggestive of the occurence of real jump in the pulsar rotation. Error bars are 1-$\\sigma$ formal standard errors.} \\end{figure}  \\subsection{Manual search for microglitches} As usual, the manual technique for microglitch detection relies largely on visual inspections of plots of carefully calculated residuals of pulsar spin frequency ($\\Delta{\\nu}$) and/or its first derivative ($\\Delta{\\dot{\\nu}}$) for scatters significantly in excess of the measurement uncertainties. The frequency and spin-down rate residuals can be calculated either directly, using coefficients of local fits of 2nd-order polynomial model to short spans of BTOAs (Flanagan 1995; Cordes et al.\\ 1988), or indirectly, by numerically differentiating phase residuals (e.g.\\ DA95; CD85). The apparent high quality of the current data (high density of BTOAs) allowed us to implement the former method, which promises to be  more sensitive to small discrete discontinuities.  For each of the 20 pulsars, a time series of the pulsar rotation frequency ($\\nu(t)$) were calculated by performing weighted least-square fits of a 2nd-order polynomial to short independent blocks of BTOAs. The block lengths were determined by the data sampling rates, the intrinsic scatter in the BTOAs and our desire to keep the formal errors in $\\nu$ to $\\sim$ 1 part in 10$^{9}$. These constraints resulted in segment lengths in the range of 35 $-$ 120 days. The frequency residuals ($\\Delta{\\nu}(t)$) were calculated as the differences between the raw data ($\\nu(t)$: obtained from local fits to independent shorts segments of BTOAs) and a model[$\\nu$, $\\dot{\\nu}$] fitted to the entire data span length (with epoch near the midpoint of the data). The frequency residuals of 4, out of the 20, pulsars used in current analysis are shown in Fig.~2. The $\\Delta{\\nu} - t$ plots were carefully inspected for evidence for possible discontinuous rotation, usually depicted by scatters whose amplitudes are significantly ($>$ 1-$\\sigma$) larger than the error bars. Generally, such large rotational discontinuities were interpreted as signatures of microglitch events and the unrefined epochs of the events were read directly from the horizontal axis of the graph. The unrefined epochs of some of the candidate microglitch events are marked with arrows in Fig.~2.  Once a candidate event had been identified in this manner, short data segments spanning $\\sim$ 200$-$400-d, and bracketing the candidate microglitch event were modelled with a 2nd-order polynomial. The resulting phase residuals were carefully re-examined for evidence for a discontinuous rotation, easily seen as a sudden slope change in the observed residuals. Such a sudden slope change is attributed to a microglitch event and the epoch is now obtained directly from the phase residual plots. Fig.~3 shows typical examples of the phase residuals from short segments known to habour one or more candidate events, first identified with the $\\Delta{\\nu}$ plots. {\\bf Basically, possible candidate events are simply depicted by sudden slope changes in the observed phase residuals} In principle, this approach, generally, results in improved epochs of candidates events, with epoch uncertainties reduced to observation intervals.  The relevant jump parameters (the epoch, the sizes of the jumps in $\\nu$ and $\\dot{\\nu}$: $\\Delta{\\nu}$ and $\\Delta{\\dot{\\nu}}$, respectively) were estimated by separately modelling the BTOA data extending between 200 and 400 d on both sides of an identified event with equation~(1). This method resulted in two sets of ephemerides for each event. $\\Delta{\\nu}$ and $\\Delta{\\dot{\\nu}}$ were estimated, respectively, from $\\Delta{\\nu} = \\nu_{\\mathrm{post}} - \\nu_{\\mathrm{pre}}$ and $\\Delta{\\dot{\\nu}} = \\dot{\\nu}_{\\mathrm{post}} - \\dot{\\nu}_{\\mathrm{pre}}$ (Chukwude 2002). The parameters $\\nu_{\\mathrm{pre}}$ and $\\dot{\\nu}_{\\mathrm{pre}}$ refer to the pre-event values extrapolated to the event epoch, while $\\nu_{\\mathrm{post}}$ and $\\dot{\\nu}_{\\mathrm{post}}$ are the corresponding post-event values, also extrapolated to the event epoch. Current method identified 124 candidate discrete jumps in $\\nu$ and/or $\\dot{\\nu}$: about 80 were noticeable simultaneously in $\\nu$ and $\\dot{\\nu}$ while 25 and 19 were identified in only $\\nu$ and $\\dot{\\nu}$, respectively.  \\subsection{Automated search for  microglitches} The current automated (or at least semi-automated) method for detecting microglitches was developed and applied in analysis of Vela pulsar data (Flanagan 1995). The technique employs a statistic, the ``goodness of fit\" parameter ($Q$), to determine when the model fit to BTOAs suddenly becomes too bad, probably, due to discrete rotational discontinuity. In HartRAO, the least-squares fitting routine employed in pulsar timing analysis uses a version of CURFIT of Bevington (1969), as modified by Flanagan (1995). Some of the outputs of the modified CURFIT include the chi-square ($\\chi^{2}$) and the reduced chi-square ($\\chi_{R}^{2} = \\chi^{2}/{dof}$, where $dof$ is the number of degrees of freedom of the fit). The goodness of fit parameter $Q$ is estimated by calculating the probability of obtaining the measured value of $\\chi_{R}^{2}$ from the data, assuming that the fitted model is the correct one. Following Flanagan (1995), $Q = Q(dof/2, \\chi^{2}/2)$, where $Q = Q(dof/2, \\chi^{2}/2)$ is the incomplete gamma function. However, the calculation of meaningful $Q$ values is strongly dependent on the availability of sensible estimates of the BTOA errors employed in the weighted least-squares fits. As a result, extra efforts are required in order to estimate realistic pulse times of arrival uncertainties provided to the fitting routine. This is extremely important in order to minimise the sensitivity of $Q$ values to seeming outlier BTOA data.  In this method, weighted least-square fits are performed. About 50$-$100-d block of BTOAs was modelled with equation~(1) plus a DM term. The block length was subsequently incremented in steps of 7 d, correponding roughly to the mean observation interval, until $Q$ value suddenly dropped below 0.1. A broad-brush approach is to interprete such a low $Q$ value as suggesting a poor description of the data by the model, due to a small discrete jump in the pulsar rotation. The resulting ephemerides were written to a file. The procedure was repeated, now starting from the epoch of this presumed event, until the end of the data set was reached (hereafter referred to as forward pass). The entire process $-$ identification of candidate events, writing of the ephemerides to a file and restarting the model-fits $-$ is automated. Hence, once the parameters are properly set up, a complete pass through the data was achieved, with minimum manual intervention. To ensure that very small amplitude events were not averaged out, the $\\sim$ 50-d start block was staggered by $\\sim$ 20 d, at a time, and the whole process was repeated. {\\bf The staggered fits improved the sensitivity of the automated search technique, as it identified a few more microglitch events which were completely averaged out in the relatively long start blocks.} Finally, the process was repeated starting from the end of a data set, and going backwards until the start of the data is reached (hereafter referred to as backward pass).  The results of these procedures were sets of ephemeris files containing ephemerides extending on both sides of the candidate microglitch events. The epoch of an event was taken as the average of the two arrival times bracketing the event and the jump parameters, $\\Delta{\\nu} = \\nu_{\\mathrm{post}} - \\nu_{\\mathrm{pre}}$ and $\\Delta{\\dot{\\nu}} = \\dot{\\nu}_{\\mathrm{post}} - \\dot{\\nu}_{\\mathrm{pre}}$, were calculated as the difference between  consecutive ephemerides after extrapolating to the event epoch. In most cases, real discrete discontinuities identified in forward pass were also detected in the backward pass with nearly coincident epochs and jump parameters of opposite signs, but of almost equal magnitudes. The automated technique identified 330 candidate discrete jumps simultaneously in $\\nu$ and $\\dot{\\nu}$. Additional 80 and 130 events were resolved separately in $\\nu$ and $\\dot{\\nu}$, respectively. As expected, the automated technique equally identified all the manually observed candidate events, as evidenced by the nearly coincident epochs and jump parameters of similar size and signs. However, it was observed that the parameters of these microglitches were generally better resolved by the manual technique. {\\bf Currently, we are not aware of any serious systematic bias in the techniques implemented in the current analysis. A possible bias in the automated search method could have resulted from BTOA data points with grossly underestimated errors. However, the resulting spurious events, which would be characterised by significantly different epochs in the forward and backward passes through the data, are easily eliminated during manual inspections of the microglitch data.}  \\begin{figure*} \\begin{minipage}{120mm} \\includegraphics[width=100mm, angle=270, clip=']{p1fig3.eps} \\caption{The phase residuals around some candidate microglitch events in the pulsar B0740--28, whose observed timing activity show strong evidence for enhanced microglitch activity. The arrows indicate epochs sudden slope changes caused discontinuous rotation} \\end{minipage} \\end{figure*}  \\subsection{Significance of the microglitch parameters} The amplitudes of all identified candidate events in $\\nu$ and/or $\\dot{\\nu}$ were tested against the null hypothesis that they are mere accumulations of random walks of preferred signs. Following  D'Alessandro et al.\\ (1995), Cordes \\& Downs (1985) the significance test was accomplished by comparing the amplitude of the observed events in $\\nu$ and $\\dot{\\nu}$ with the expected standard deviations of idealized large-rate random walk process. {\\bf This refers to an accumulation of step functions in pulsar phase residuals and/or its derivatives whose effects, over a time, could mimic genuine discrete rotational discontinuities.} Events are considered real, that is their amplitudes are too large to be due to random walk fluctuations, if $\\Delta{\\nu} \\geq \\sigma_{\\Delta{\\nu}}N$ and $\\Delta{\\dot{\\nu}} \\geq \\sigma_{\\Delta{\\dot{\\nu}}}N$: where $\\sigma_{\\Delta{\\nu}} = \\sqrt{S_{1}\\Delta{t}}$ and $\\sigma_{\\Delta{\\dot{\\nu}}} = \\sqrt{S_{2}\\Delta{t}}$ are, respectively, the standard deviations of idealized random walks in $\\nu$ and $\\dot{\\nu}$ with strength parameters $S_{1}$ and $S_{2}$, respectively, $N$ is an integer denoting the significance level of the test and $\\Delta{t}$ is the difference between the arrival times bracketing a candidate event. In this paper, $\\Delta{t}$ represents the upper bound on the rise time of the microglitch events. For the current sample, values of $\\Delta{t}$ lie between $\\sim$ 0.05 and 40 d, with a mean value of 10 d.  Several authors (e.g.\\ CD85, Cordes et al.\\ 1988; DA95) have demonstrated that $N = 5$ is an appropriate signficant level, as it correctly rejects spurious random walk-induced events which could mimic real microglitches. Given the similarity between the HartRAO and JPL pulsar data, in terms of observing frequencies and BTOA signal-to-noise ratio,  we equally adopted $N = 5$ for the current significance test.  Following Chukwude (2002), DA95, CD85, we calculated strength parameters ($S_1$ and $S_2$, respectively) from BTOAs bracketed by  any two identified candidate events (hereafter referred to as method A) and from the entire BTOA data available for a pulsar (hereafter referred to as method B). Our presumption is that data sets used in method A, probably, contain no discrete discontinuity and hence has extremely low timing activity. On the other hand, the data set employed in method B incorporates all the rotational discontinuities, including all the microglitches. In principle, methods A and B are expected to yield, respectively, the lower and upper limiting estimates for the strength parameters. The bottom line is that, though the strength parameter estimate is variable in the presence of random walks, its true value should lie within the values obtained from the two methods. Since the rms phase residuals of short and long data span lengths are believed to be dominated by events in $\\nu$ and $\\dot{\\nu}$, respectively (DA95; CD85). In view of this fact, we placed greater weight on methods A and B in testing the significance of $\\Delta{\\nu}$ and $\\Delta{\\dot{\\nu}}$, respectively. In addition, the amplitudes of a candidate events were required to exceed the formal standard error by, at least, a factor of two. In principle, the necessary condition that must be satisfied by the jumps in $\\nu$ and $\\dot{\\nu}$ is that $\\Delta{\\nu} > 2\\epsilon_{\\Delta{\\nu}}$ and $\\Delta{\\dot{\\nu}} > 2\\epsilon_{\\Delta{\\dot{\\nu}}}$, respectively, where $\\epsilon_{\\Delta{\\nu}}$ and $\\epsilon_{\\Delta{\\dot{\\nu}}}$ are, the formal standard errors in $\\Delta{\\nu}$ and $\\Delta{\\dot{\\nu}}$, respectively. The formal standard errors were obtained directly from the HartRAO in-house timing analysis routine. We believe that these conditions are stringent enough to filter out most, if not all, spurious events, which could mimic genuine microglitch events.  Over 600 candidate discrete jumps were identified by a combination of the manual and automated search techniques for microglitches. However, only 299 of these events have amplitudes that are large enough to pass the significance tests (and are hereafter referred to as microglitches). The remainder of $\\sim$ 300 jumps failed the tests and were consequently rejected, as mere spurious events arising from accumulations of large rate steps of random walks. Out of the 299 microglitches, 266 ($\\sim$ 90\\%) are found to be simultaneously significant in both $\\nu$ and $\\dot{\\nu}$ ( 71 of these were better resolved by the manual search technique). {\\bf These represent microglitch events that produced detectable changes in the pulsar spin rate evolution. However, it is believed that several candidate events might have been averaged out in the process leading to the calculation of the $\\Delta\\dot{\\nu}$ data. Hence, one expects the more objective automatic technique to resolve several more microglitches.} Others are noticeable in either of the spin parameters ($\\nu$ or $\\dot{\\nu}$). The parameters of the 266 and 33 microglitches are summarized in Tables~(1) and (2), respectively. A summary of our results shows that the total microglitches in $\\nu$ and $\\dot{\\nu}$ are 285 and 280, respectively. Of the 285 microglitches in $\\nu$, 148 and 137 of the events have positive and  negative amplitudes, respectively, while 139 and 141 of $\\Delta{\\dot{\\nu}}$ are, respectively, positive and negative. The fractional jumps in $\\nu$ ($\\Delta{\\nu}/\\nu$) have minimum and maximum values as $\\sim$ $-8\\times10^{-9}$ and $2\\times10^{-8}$, respectively, while $\\sim$ $-3\\times10^{-3}$ and $4\\times10^{-3}$ correspond, respectively, to minimum and maximum amplitudes of $\\Delta{\\dot{\\nu}}/\\dot{\\nu}$. The jumps, nonetheless, show no apparent preferential combination of signs ($\\Delta{\\nu}$, $\\Delta{\\dot{\\nu}}$).  \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Parameters of the microglitches, whose amplitudes are simultaneously significant in both $\\nu$ and $\\dot{\\nu}$, for the 20 HartRAO radio pulsars observed using both the manual and automatic search techniques.} \\begin{tabular}{@{}llrrrrcllrrrr} \\hline \\multicolumn{1}{l}{Pulsars} & \\multicolumn{1}{l}{Epoch} & \\multicolumn{1}{r}{$\\Delta{\\nu}/\\nu$} & \\multicolumn{1}{r}{ERR$^{a}$} & \\multicolumn{1}{r}{$\\Delta{\\dot{\\nu}}/\\dot{\\nu}$} &  \\multicolumn{1}{r}{ERR$^{a}$} &  &   \\multicolumn{1}{l}{Pulsars} & \\multicolumn{1}{l}{Epoch} & \\multicolumn{1}{r}{$\\Delta{\\nu}/\\nu$} & \\multicolumn{1}{r}{ERR$^{a}$} & \\multicolumn{1}{r}{$\\Delta{\\dot{\\nu}}/\\dot{\\nu}$} &  \\multicolumn{1}{r}{ERR$^{a}$} \\\\ (PSR B) & (MJD) & ($10^{-9}$) &   & ($10^{-3}$) &  &  & (PSR B) &  (MJD) & ($10^{-9}$) &   & ($10^{-3}$) &  \\\\ (1)  & (2)  & (3) & (4) & (5) & (6) &  & (1)  & (2)  & (3) & (4) & (5) & (6) \\\\\\hline 0736$-$40&47076&0.43&3&$-$0.15&2& &0835$-$41&46271&$-$0.43&5&2.2&4 \\\\ & 47751  &  0.12 &2&$-$0.16&2 & & & 47207  &  0.07&1&0.11&1  \\\\ & 47943$^{+}$ &  0.25&3&0.17&2 & & & 47840  &  0.12&1&$-$0.079&3\\\\ & 48283  &  0.13&2& 0.75&2& & &49780  & $-$0.13&6&$-$0.064&3\\\\ & 48502  & 1.08 & 8 &$-$0.89 & 4& & 0959$-$54& 46479 & 6&2&$-$1.8&5\\\\ & 49292$^{+}$ &  0.34&4&0.83 &6& & & 47477$^{+}$ &  2.2&2& 0.32 &6 \\\\ & 50395$^{+}$ & $-$0.67&2&$-$0.44&2& & & 48138$^{+}$ &  2.4&4&$-$0.64 &2\\\\ & 50612  & 0.17&2&0.42&1 & & & 48490  & $-$1.64&8&$-$0.46&2    \\\\ & 51997  & 0.13&2&$-$0.83&3& &  & 48600 &  0.8&2& 0.75 &9\\\\ & 52224  & 0.45&6&0.55&8 & & & 49241$^{+}$ & $-$1.79 &8& 0.80 &6 \\\\\\\\ & 52409$^{+}$& $-$0.86&5&$-$1.2&4& & & 50845$^{+}$ &  1.74&8& 1.19 &4\\\\ 0740$-$28&46653$^{+}$&0.46&7&$-$0.052&5& & &52236$^{+}$ & $-$3.7 & 1& 1.98 & 9 \\\\ & 46806  &  1.2&4&$-$0.08&2& & 1054$-$62 & 47153 & 0.37&8&$-$0.28&5 \\\\ & 46992  & $-$0.17&6&0.057&6& & & 47913  &  0.35&5&0.19&3 \\\\ & 47085  & $-$0.06&1&$-$0.026&5& & &48190$^{+}$ &$-$0.95&8& 0.26 &6 \\\\ & 47358$^{+}$ & 0.72&3&$-$1.05&9& & & 48301  &  0.69&6&0.18&3 \\\\ & 47450  & 0.39&5&$-$0.146&7& & & 48480$^{+}$ & 1.5&2&$-$0.56 &3\\\\ & 47576  & $-$0.38&8&$-$0.08&2& &  & 49169  & 0.12&4&0.06&2 \\\\ & 47681$^{+}$ & 2.75&8&$-$1.6 & 2& &  & 49378  &  0.29&8&$-$0.19&6\\\\ & 47736  & $-$0.55&9&0.09&2& & & 49631  & 0.18&5&$-$0.09&2 \\\\\\\\ & 47843  & $-$0.75&5&$-$0.10&1 & & & 50025  & $-$0.5&2&0.12&3  \\\\ & 47956  & $-$0.95&6&0.070&5 & & & 50386  & 0.20&5&0.47&2\\\\ & 48084  & $-$0.62&7&0.034&6 & & & 50782$^{+}$ &  3.2&4&$-$0.88&8 \\\\ & 48246$^{+}$ & $-$2.3&3&$-$0.99 &2 & & & 50997$^{+}$ &  2.3&4&$-$0.59&5 \\\\ & 48345  &  1.4&2&2.5 & 4 & & & 52008  & $-$0.12&3&0.12&3 \\\\ & 48427  & $-$0.7 & 2 &$-$0.29 & 4 & & & 52166  & $-$0.75&9&$-$0.11&5 \\\\ & 48470$^{+}$&$-$2.4&2& 2.6 &2& & 1221$-$63& 48147 & $-$0.06 &1&$-$0.022&1 \\\\ & 48646  &  1.14 & 9 & $-$0.07 &1 & & & 48788$^{+}$ &  0.87&4&0.79 & 7\\\\ & 48852$^{+}$  &  1.23  & 8 &0.25&2& & & 49308  &  0.06&2&0.05&1 \\\\ & 48940  &  1.2 &2&0.6&2 & & & 49814$^{+}$ &  1.5 &4&$-$0.19&6 \\\\\\\\ & 49050  &  2.8&3 & 1.55 &3 & & & 50830  & $-$0.17&4&0.04&1 \\\\ & 49287  &  2.0 &2 & 0.96&3 & & & 50960$^{+}$ & 3.6 &5& 0.96&9 \\\\ & 49353$^{+}$ & $-$2.2 & 5 & 0.35 &2 & & & 51931  &  1.2&3&$-$0.9&3 \\\\ & 49621& $-$0.16&4&$-$0.092&4& & 1240$-$64& 47460 & 0.18&2&$-$0.12&3\\\\ & 49915  &  0.89&4&0.070&3 & & & 47960$^{+}$ & $-$0.33&4&0.09&2 \\\\ & 50008  & $-$0.5&2&0.16&2 & & & 48464  & 0.19&2& 0.12&3 \\\\ & 50185 &  0.65&8&$-$1.3&2 & & & 49030  & $-$0.033&2& 0.29&2 \\\\ & 50376$^{+}$ & 1.19&9& $-$0.9&2 & & & 49610  &  0.32&3&$-$0.13&1 \\\\ & 50529  & 0.77&8&$-$0.08&1& & & 50092  & $-$0.68&5&0.18&2 \\\\ & 50609&0.4&1&$-$0.18&2& & &50503$^{+}$&$-$0.78&6&$-$0.59&9\\\\\\\\ & 50782  & 1.3&2&1.0&2 & & & 50677  & $-$0.11&3&0.014&2 \\\\ & 50978$^{+}$&1.45&8&0.69 &3& & & 51871& $-$0.09&2&$-$0.06&1 \\\\ & 51109  & 1.5&2&$-$0.53&3 & & 1323$-$58 & 47586$^{+}$ & 1.43&8&1.5&3\\\\ & 51785$^{+}$ & $-$1.93&8&0.96&2& & & 47878  &  $-$0.16&2&1.30&2 \\\\ & 51880  & $-$ 0.93 & 8 & 0.38 & 2& & &47955$^{+}$ & $-$1.45&5&1.07&8\\\\ & 51970  & $-$0.19&7&$-$0.22&1& & & 48212  & $-$1.03&4&0.85&2\\\\ & 52030$^{+}$ & 2.32&9&$-$1.38&9 & & & 48700  & $-$0.69&2&$-$1.28&1 \\\\ & 52111  & $-$0.7&2&$-$0.11&3 & & & 49245  & $-$0.78&5&0.43&2 \\\\ & 52249  & $-$1.2&2&0.08&2 & & & 49792  & $-$1.82&7&$-$0.82&3\\\\ & 52370$^{+}$ & $-$0.28&5&0.25 &2 & & & 49879  &  0.47&4&0.28&2\\\\\\hline \\multicolumn{13}{l}{$^{a}$ The formal standard error and refers to the least significant digit}   \\\\ \\multicolumn{13}{l}{$^{+}$ Indicates epochs of events whose jump parameters were better resolved by the manual technique} \\\\\\hline  \\end{tabular} \\end{minipage} \\end{table*}  \\begin{table*} \\centering \\begin{minipage}{140mm} \\contcaption{} \\begin{tabular}{@{}llrrrrcllrrrr} \\hline \\multicolumn{1}{l}{Pulsars} & \\multicolumn{1}{l}{Epoch} & \\multicolumn{1}{r}{$\\Delta{\\nu}/\\nu$} & \\multicolumn{1}{r}{ERR$^{a}$} & \\multicolumn{1}{r}{$\\Delta{\\dot{\\nu}}/\\dot{\\nu}$} &  \\multicolumn{1}{r}{ERR$^{a}$} &  &  \\multicolumn{1}{l}{Pulsars} & \\multicolumn{1}{l}{Epoch} & \\multicolumn{1}{r}{$\\Delta{\\nu}/\\nu$} & \\multicolumn{1}{r}{ERR$^{a}$} & \\multicolumn{1}{r}{$\\Delta{\\dot{\\nu}}/\\dot{\\nu}$} &  \\multicolumn{1}{r}{ERR$^{a}$} \\\\ (PSR B) &  (MJD) & ($10^{-9}$) &   & ($10^{-3}$) &  &  & (PSR B) &  (MJD) & ($10^{-9}$) &   & ($10^{-3}$) &  \\\\ (1)  & (2)  & (3) & (4) & (5) & (6) &  & (1)  & (2)  & (3) & (4) & (5) & (6) \\\\\\hline & 50150$^{+}$ &  0.99&9&1.23&9& & & 47860 & 1.5 &4 &$-$1.9&5\\\\ & 50455  & $-$0.92&6&0.98&3 & & & 48000  &   1.1&3&$-$1.2&3 \\\\ & 50670$^{+}$ & $-$0.55&8&$-$0.58&8 & & & 48316& $-$0.99&9&$-$0.32&2 \\\\ & 50964 & 0.48 & 8 & $-$0.64 & 5& & & 48548  & 0.75 &6& 0.54 & 5 \\\\ & 51020  &  0.5& 1 &$-$0.56&5& & & 49085$^{+}$ &$-$1.9&2&$-$2.46&1 \\\\ & 51880  & 1.08 & 7 & $-$0.66 & 4 & & & 49206  &  0.60&3&1.78&9 \\\\ & 52338  &$-$1.5 & 3 & $-$1.2 & 4& & & 50210& $-$1.5 & 2 &$-$0.57&3 \\\\ 1323$-$62 & 46769 & $-$0.7&2&0.06&2& & &50673$^{+}$&$-$1.7 &8&$-$2.58&6\\\\ & 46940  & $-$1.92&10&$-$0.75&1& & & 52276  &  1.6& 2&1.52&10 \\\\ & 47163&$-$1.4 & 3&0.28&4 & & 1449$-$64& 48360&0.22&1&0.042&6 \\\\\\\\ & 47666$^{+}$ &  0.82&6&$-$0.17&1& & & 50212  &   0.05&1&0.014&4 \\\\ & 47729$^{+}$ &  $-$0.90&5&0.8&2& & & 50768  &   0.12&5&0.05&2 \\\\ & 47875  &   3.8 & 8 &0.6&2 & & & 50835  &  $-$0.08&3&$-$0.06&1 \\\\ & 48059  &   0.40&7&$-$0.19&4& & 1556$-$44&48861&$-$0.06&1&0.008&3\\\\ & 48242$^{+}$ &  $-$1.3& 4 &0.8&2 & & & 50104  &   0.08&1&$-$0.05&1 \\\\ & 48371  &  $-$0.63&8&0.32&3 & & & 51166  &   0.15&5&$-$0.08&1 \\\\ & 48459$^{+}$ &   0.67&9&0.48&5& &1557$-$50&47413&$-$0.12&2&$-$0.035&5 \\\\ & 48670  & $-$0.56&6&$-$0.9&2& & & 47913  &  $-$0.04&1&0.0130&2 \\\\ & 48860$^{+}$ &  0.69&9&0.65&4& & & 47960  &  $-$0.18&2&$-$0.025&4 \\\\ & 48939  & $-$1.4 & 2 &0.55&4 & & &48392$^{+}$&$-$0.66&5&$-$0.095&3 \\\\\\\\ & 49150$^{+}$ &  1.6& 2 &0.19&3 & & & 48524  &   0.12&2&0.058&4\\\\ & 49334  & $-$1.5 & 3 &$-$0.38&4& & & 48552 &  $-$0.18&4&$-$0.083&6 \\\\ & 49408  & $-$0.62&8&$-$0.44&3 & & & 48928$^{+}$ & 0.74&3&$-$0.19&4 \\\\ & 49581  &  0.9 & 2 &0.16&4 & & & 49326  &  $-$0.12 &2&0.06 &1 \\\\ & 49812  &  0.42&8&$-$0.26&2 & & &49750$^{+}$ & $-$0.28&3&0.092&5  \\\\ & 49906  &  $-$0.36&5&0.07&2 & & & 50018  &  $-$0.06&1&0.10&1 \\\\ & 50019  &   1.3 & 2 &0.18&3 & & & 50764$^{+}$ &  $-$0.52&2&0.076&5 \\\\ & 50104  & $-$0.9& 2 &$-$0.20&4 & & & 51274  &  $-$0.32&5&$-$0.07&2 \\\\ & 50255$^{+}$ &  $-$1.27&9&$-$0.38&5 & & & 52007 &$-$0.09&3&0.06&2 \\\\ & 50351  &   2.2& 3&$-$0.64&9 & & 1642$-$03 & 47156 & 2.15&7&1.93&4 \\\\\\\\ & 50414  & $-$1.07& 9&0.23&4 & & & 47740& 0.45&3&$-$1.28&3 \\\\ & 50646 &  $-$2.6&3&$-$0.88&6 & & & 48164$^{+}$ &  $-$2.81&3&1.79&6 \\\\ & 50824  &  $-$1.3 &3&0.41&8 & & & 48939  &   0.33&2&$-$0.51&1\\\\ & 50970$^{+}$ &  4.5&8&1.7&5 & & & 49358  &   1.16& 9 &0.70&6 \\\\ & 51167  &   1.6&2&$-$0.24&8 & & & 49646  &  $-$0.42&2&$-$0.85&2\\\\ & 51950 &  $-$1.8&4&$-$0.5&1 & & & 50280  &  $-$0.59&3&1.24&3\\\\ & 52073  &  $-$2.1&4&0.39&6& & & 50650$^{+}$ & 1.9&2&1.48&10 \\\\ & 52264$^{+}$ &   1.9 &2&0.67&4 & & & 50978  &  $-$0.22&3&$-$1.59&2 \\\\ 1356$-$60 & 47855 & 0.21&4&$-$0.12&1& & & 51173&$-$0.8 & 2 &$-$1.3&2 \\\\ & 48680  & $-$1.9&5&$-$0.9 &2 & & & 52183$^{+}$ & $-$1.85&8&1.72&2 \\\\\\\\ & 48847& 0.45&7&$-$0.033&8 & & 1706$-$16&48131$^{+}$&$-$3.5&8&2.7&5\\\\ & 49388  &  0.38&6&$-$0.06&2 & & & 48535 &  $-$0.89& 9 & $-$1.7&4 \\\\ & 49685  &   0.65&4&0.047&3 & & & 48615  & 0.56 & 5 & 3.3 &3 \\\\ & 49713  &  $-$0.53&8&$-$0.04&1 & & & 49101$^{+}$ & $-$3.6 &5& $-$3.5&6 \\\\ & 50101$^{+}$ & $-$1.68&9&$-$0.85&1 & & & 49257  &  1.15&8&$-$0.65&3 \\\\ & 50356  &  0.022&4&0.049&8 & & & 50541  &  1.0&2&0.31&5 \\\\ & 50757  &  0.26& 6 &$-$0.056&6& & &51002& 0.30&6&$-$0.22&1\\\\ & 51150  &   0.8& 2 &0.85&3& & & 51830 &$-$0.48&4&$-$0.36&6 \\\\ & 52009 &$-$0.12&4&$-$0.07&1& & 1727$-$47&46383&$-$1.5 &5 &0.06&2 \\\\ 1358$-$63&47725$^{+}$&$-$1.8& 3&$-$2.38 & 9& & &46913&$-$1.2&4&0.034&8 \\\\\\hline \\multicolumn{13}{l}{$^{a}$ The formal standard error and refers to the least significant digit}   \\\\ \\multicolumn{13}{l}{$^{+}$ Indicates epochs of events whose jump parameters were better resolved by the manual technique} \\\\\\hline \\end{tabular} \\end{minipage} \\end{table*}  \\begin{table*} \\centering \\begin{minipage}{140mm} \\contcaption{} \\begin{tabular}{@{}llrrrrcllrrrr} \\hline \\multicolumn{1}{l}{Pulsars} & \\multicolumn{1}{l}{Epoch} & \\multicolumn{1}{r}{$\\Delta{\\nu}/\\nu$} & \\multicolumn{1}{r}{ERR$^{a}$} & \\multicolumn{1}{r}{$\\Delta{\\dot{\\nu}}/\\dot{\\nu}$} &  \\multicolumn{1}{r}{ERR$^{a}$} &  & \\multicolumn{1}{l}{Pulsars} & \\multicolumn{1}{l}{Epoch} & \\multicolumn{1}{r}{$\\Delta{\\nu}/\\nu$} & \\multicolumn{1}{r}{ERR$^{a}$} & \\multicolumn{1}{r}{$\\Delta{\\dot{\\nu}}/\\dot{\\nu}$} &  \\multicolumn{1}{r}{ERR$^{a}$} \\\\ (PSR B) &  (MJD) & ($10^{-9}$)  &   & ($10^{-3}$) &  &  & (PSR B) &  (MJD) & ($10^{-9}$) &   & ($10^{-3}$) &  \\\\ (1)  & (2)  & (3) & (4) & (5) & (6) &  & (1)  & (2)  & (3) & (4) & (5) & (6) \\\\\\hline & 47123& 0.5 & 1 &$-$0.017&2& & & 46979  &   0.92&9&0.42&8 \\\\ & 47260$^{+}$&9.3 & 3&$-$1.95&6& & &47132& 1.2&6&$-$0.48&9 \\\\ & 47499 &$-$0.33&4&$-$0.07&2& & & 47322$^{+}$ & 3.5 &6& 0.75&8\\\\ & 47889  &   0.6 & 2 &0.50&5& & & 47764  &   2.4&7&0.3&1 \\\\ & 48540  &  $-$0.26&8&0.08&1& & & 48140$^{+}$ & 1.9&5& 0.25&5\\\\ & 49200$^{+}$&13.6 &8&$-$2.54 &2 & & &48210 &$-$2.07&8&0.5&2\\\\ & 49817 &   0.84&7&$-$0.12&1 & & &48451  &   0.6&2&0.08&2 \\\\ & 50088  &$-$1.13&6&$-$1.04&2& & & 48580 & $-$7.8&2&$-$0.89&9 \\\\ & 50388&0.97&7&0.033&1 & & & 48747  &  $-$0.9&4&$-$0.16 &4\\\\ & 50662& 1.3 &3&0.38&8 & & &49132&$-$1.7&6&$-$0.7& 3 \\\\\\\\ & 52108 & 0.7&2&0.05&1 & & & 49279  &  $-$4.7&9&0.6& 2 \\\\ & 52300 & 3.76 & 8 & $-$0.58& 6 & & & 49876& 3.6&2&$-$0.38&2 \\\\ 1749$-$28 & 46805 & $-$0.50&4&0.12&1& & & 50150$^{+}$& 4.9 & 3& $-$2.35& 7 \\\\ & 46990$^{+}$& 5.03 & 7 & $-$1.6 & 2& & & 50495  &  $-$1.2&2&0.18&6 \\\\ & 47315 &  0.58&2&0.23&2 & & & 50585  & $-$0.8&2&$-$0.28&3 \\\\ & 47660 & $-$0.31&2&$-$0.13&1 & & & 50997  &$-$0.98&4&$-$0.68&9\\\\ & 48098$^{+}$ & 3.5 &8&0.79&3 & & & 51905$^{+}$ & 32 &3& 1.68&8 \\\\ & 48485 &  0.48&6&$-$0.42&4 & & & 52041  & $-$0.83&4&$-$0.74&9 \\\\ & 48701$^{+}$ & $-$1.3 & 2&0.89 &6 & & & 52110 & 0.8&3 &1.92&8  \\\\ & 48806 & $-$0.96&6&0.36&5 & &  & 52273  &  12.6&6  & $-$2.22&5 \\\\\\\\ & 48916$^{+}$ &2.3&5&0.62&9 & &1929+10&47391&$-$0.027&6&$-$0.067&8\\\\ & 49173 & $-$0.32&6&$-$0.18&5 & & &47670$^{+}$&$-$0.62&4 &$-$0.21 &8\\\\ & 49756$^{+}$ &  0.8&2&$-$0.6 & 1& & & 48360$^{+}$ & 0.74&7&$-$0.39&2\\\\ & 49964 &  0.30&5&0.57&3 & & & 48446  &  $-$0.28 & 2&0.06&1 \\\\ & 50260 &  1.09&6&0.41&3 & & & 49064  &  $-$0.23 & 4&$-$0.16&2 \\\\ & 50356 & $-$0.33&5&$-$0.39&4 & & & 49205  &  $-$0.27 & 5&0.15&2\\\\ & 50408$^{+}$ &  0.56&6&$-$1.12&8 & & &49542& $-$0.16 & 5&0.18&3\\\\ & 50583 & $-$0.31&6&0.13&4 & & &49710&$-$0.013 & 2&$-$0.012&2 \\\\ & 50727 &  0.68&9&$-$0.15&2& & & 50182  &  $-$0.05 & 1&0.33&4 \\\\ & 51108$^{+}$ & 1.2 &2&$-$0.65&7 & & &50266&$-$0.07 & 2&$-$0.38&3\\\\\\\\ & 51250 &  0.65&6&$-$0.21&1 & & & 50773  &   0.09&2 & 0.14&2\\\\ & 52075$^{+}$ & $-$0.76&4&0.19&1 & & &51198& 0.19&3 & 0.10&2 \\\\ 1822$-$09&46838& $-$2.2&5&$-$0.17&6 & & &52081&0.5&1& $-$0.25&3\\\\\\hline \\multicolumn{13}{l}{$^{a}$ The formal standard error and refers to the least significant digit}   \\\\ \\multicolumn{13}{l}{$^{+}$ Indicates epochs of events whose jump parameters were better resolved by the manual technique.} \\\\\\hline \\end{tabular} \\end{minipage} \\end{table*}   \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Parameters of the 33 microglitches, whose amplitudes are significant in only $\\nu$ or $\\dot{\\nu}$, observed using the automated search technique} \\begin{tabular}{@{}llrrrrcllrrrr} \\hline \\multicolumn{1}{l}{Pulsars} & \\multicolumn{1}{l}{Epoch} & \\multicolumn{1}{r}{$\\Delta{\\nu}/\\nu$} & \\multicolumn{1}{r}{ERR$^{a}$} & \\multicolumn{1}{r}{$\\Delta{\\dot{\\nu}}/\\dot{\\nu}$} &  \\multicolumn{1}{r}{ERR$^{a}$} &  & \\multicolumn{1}{l}{Pulsars} & \\multicolumn{1}{l}{Epoch} & \\multicolumn{1}{r}{$\\Delta{\\nu}/\\nu$} & \\multicolumn{1}{r}{ERR$^{a}$} & \\multicolumn{1}{r}{$\\Delta{\\dot{\\nu}}/\\dot{\\nu}$} &  \\multicolumn{1}{r}{ERR$^{a}$} \\\\ (PSR B) &  (MJD) & ($10^{-9}$)  &   & ($10^{-3}$) &  &  & (PSR B) &  (MJD) & ($10^{-9}$) &   & ($10^{-3}$) &  \\\\ (1)  & (2)  & (3) & (4) & (5) & (6) &  & (1)  & (2)  & (3) & (4) & (5) & (6) \\\\\\hline 0835$-$41&47437& $-$ &$-$ & $-$0.099 & 3& & & 50607 & $-$1.53 & 5 & $-$ & $-$ \\\\ & 50961&$-$0.25 & 1 & $-$ & $-$& &1240$-$64 & 47423 & 0.26 & 2 & $-$&$-$ \\\\ & 51201 & $-$0.36 & 3 & $-$ & $-$ & &  & 48712 & $-$ & $-$ & $-$0.18  & 5 \\\\ 0959$-$54 & 47570 & $-$0.95&4& $-$ &$-$& & & 49076 &$-$&$-$& $-$0.65 & 2\\\\ & 47983 & $-$2.37&4&$-$&$-$& &1323$-$62 & 49688 & $-$ & $-$ &  0.18&3 \\\\ & 48017 & 1.55&5& $-$& $-$ & &1358$-$63 & 49733 & $-$ & $-$ &$-$0.47 & 2  \\\\ & 49076 & $-$ & $-$& $-$0.65&2 & & & 48555 & $-$ & $-$ & 0.69 & 3  \\\\ & 49511 & $-$1.8 & 3 &$-$ & $-$ & & & 49767 & $-$ & $-$ & 0.49 & 6  \\\\ & 50032 & 1.82& 6&$-$& $-$& & & 50251 & $-$ & $-$ & $-$0.69  & 8  \\\\ & 50455 & 1.29 &4 &$-$ &$-$& &1557$-$50 & 47599 & 1.29 & 9 & $-$ & $-$\\\\\\\\ & 49711 & $-$ & $-$ & 0.25 & 6 & & & 49006 & $-$0.29 & 6 & $-$ & $-$ \\\\ & 50696 & $-$ & $-$ &0.5 & 1& &1749$-$28 & 50395 & $-$& $-$ & 0.79& 5\\\\ & 50983 & 0.85&3&$-$ & $-$& & 1822$-$09 & 47960 & $-$&$-$&0.12&2 \\\\ & 51899 & 3.85  & 5 & $-$ & $-$ & & & 48797 & $-$ & $-$ & $-$0.039 & 6 \\\\ & 52210 & $-$2.2 &  2 & $-$ & $-$ & & & 50098 & $-$5.7  & 7  & $-$ & $-$ \\\\ 1706$-$16 & 49880 & $-$0.86 & 8&$-$ & $-$ & & & 51908&$-$2.5&3&$-$ & $-$ \\\\ 1727$-$47 & 48886 & 0.35 & 5 & $-$ & $-$ & & & & & &  & \\\\\\hline \\multicolumn{13}{l}{$^{a}$ The formal standard error and refers to the least significant digit}.\\\\ \\\\\\hline \\end{tabular} \\end{minipage} \\end{table*} ",
        "conclusions": "\\subsection{Macroglitches among the HartRAO pulsars} Twenty-five of the twenty-six pulsars on routine monitoring at HartRAO are aged more than $10^5$ years, having an average of $10^{6.17}$ years. None of the twenty-five has shown any fractional frequency jump larger than $\\sim 10^{-8}$ within the seventeen years of observation. This probably suggests that the interval for large macroglitches in these ``middle-aged'' (characteristic age $\\sim 10^5\\ - \\ 10^7$ years) pulsars, if they do occur at all, could be more than the total observation span of $\\sim$ 425 years.  The only one in the sample that has shown a large jump, PSR 1727$-$47, is younger than $10^5$ years. Within this period, the pulsar had two  large macroglitches of amplitudes $\\Delta{\\nu /} \\nu = 1.37 \\times 10^{-7}$ and $1.25\\times 10^{-7}$ at MJD 49382 and 52476, respectively. There was also a smaller one, though larger than the jumps reported here, at MJD 50718 with $\\Delta{\\nu /} \\nu = 3.2 \\times 10^{-8}$. Very many other microglitches have been reported in this paper for this pulsar.  {\\bf This is in agreement with earlier suggestions that the oldest pulsars, like the millisecond pulsars are  free from macroglitches. In a statistical study of 48 macroglitches from 18 pulsars, Lyne et al.\\ (2000) found that the glitch activity decreases linearly with decreasing rate of slowdown. Also, Urama \\& Okeke (1999) reported that smaller values of $\\Delta{\\nu}/\\nu$ and $\\Delta{\\dot{\\nu}}/\\dot{\\nu}$ are seen to be more common in pulsars older than 10$^{5}$ years and that no glitch has been observed in conventional radio pulsars older than 10$^{7}$ years.}  \\subsection{Microglitches and timing noise events} The observed timing activity of 20 pulsars are found to be at levels (with $\\sigma_{R} / \\sigma_{W} > 5$) considered significant for meaningful microglitch analyses. For the remainder of 6 pulsars, the observed timing activity over $\\sim$ 13 yr period is almost indistinguishable from measurement uncertainty. The intrinsic scatter in the BTOAs of these objects have amplitudes in excess of 20 ms (see Chukwude 2007) which could easily and effectively swamp any low level timing activity in the pulsars. This anomalously large scatter in the data could be attributed to HartRAO local observing system parameters: high observing frequencies (13 and 18 cm) and the narrow receiver band of 10 MHz. These parameters could combine to further reduce the precision of measurements of the pulse times of arrival. This would ultimately degrade the BTOAs of pulsars which intrinsically exhibit  low level of timing activity. Another possibility is that some of the pulsars are intrinsically faint objects. For instance, previous independent analyses of, presumably, higher precision, timing data on some of these pulsars (1133+16, 1426$-$66, 1451$-$68 and 2045$-$16) do not reveal any appreciable level of timing activity (DA95; CD85).  For the sub-sample of 20 pulsars extensively used in current analyses, the results are both phenomenal and unprecedented in the history of microglitch observations in radio pulsars. The automatic technique identified a total of 266 jumps that are simultaneously significant in both pulsar rotation frequency and spin-down rate. Only 71 (or $<$ 30\\%) of these microglitches were equally resolved by the traditional manual method.  In addition, the automated technique observed 19 and 14 significant jumps in only $\\nu$ and $\\dot{\\nu}$, respectively. The fractional jumps in $\\nu$ and $\\dot{\\nu}$ ($\\Delta{\\nu}/\\nu$ and $\\Delta{\\dot{\\nu}}/\\dot{\\nu}$, respectively) have all possible combinations of signs. For the 266 microglitch events, we found ($\\Delta{\\nu}/\\nu$, $\\Delta{\\dot{\\nu}}/\\dot{\\nu}$) = (+, +), (+, $-$), ($-$, $-$) and ($-$, +) as the signatures of 69, 68, 66 and 63 events, respectively. It is worthy of note that the distribution of these events is not statistically different from random. Our results also reveal large dispersions, $\\sim$ 4 orders of magnitude, in observed sizes of the jumps in both spin parameters. Strikingly, the observed fractional jumps in $\\nu$  and $\\dot{\\nu}$ have sizes whose magnitudes are in the range of $\\sim$ 2$\\times10^{-11}$ $-$ 2$\\times10^{-8}$ and 5$\\times10^{-6}$ $-$ 4$\\times10^{-3}$, respectively. Typical median values for $|\\Delta{\\nu}/\\nu|$ and $|\\Delta{\\dot{\\nu}}/\\dot{\\nu}|$ are $\\sim$ $0.8\\times10^{-9}$ and $0.4\\times10^{-2}$, respectively.  Arguably, we have presented what is as yet the most in-depth and successful observations of microglitch phenomenon in  slow radio pulsars. Nonetheless, there had been some attempts in the past to observe these small amplitude discrete events in pulsar rotation rates. Cordes \\& Downs (1985) carried out extensive analyses of the BTOAs of 17 pulsars obtained at Jet Propulsion Laboratory (JPL) between 1968 and 1982. In the framework of the manual technique, the authors identified a total of 74 candidate events in $P$ and $\\dot{P}$. However, only 43 of these events were found to be significant $-$ 10 simultaneously in $P$ and $\\dot{P}$, while 17 and 16 were separately significant in $P$ and $\\dot{P}$, respectively. Similar analysis, based on a sample of 26 pulsars observed at Mt. Pleasant Observatory between 1987 and 1991 (D'Alessandro et al.\\ 1995) yielded about 35 significant microjump events from about 70 candidate events. A breakdown of the events shows that while 5 events were noticeable in both $\\nu$ and $\\dot{\\nu}$, 20 and 10 were significant only in $\\nu$ and $\\dot{\\nu}$, respectively. In comparison, current analyses identified $\\sim$ 300 microjumps from just 20 pulsars, which is about a factor of 4 in excess of all the hitherto known radio pulsar microglitches.  As a check on the reliability of the methods and the reality of the results, it might be necessary to compare it with those of previous studies. It is noteworthy that the HartRAO and Tasmanian timing data on 16 pulsars overlap each other, up to 1991. Remarkably, over 90\\% of the microglitch events reported for these objects by DA95 also show up in current analyses and the results (epochs and signatures of the events) are generally in good agreements. However, in terms of the jump sizes, our results differ considerably with those published by DA95, nonetheless they are still of the same order of magnitude. Irrespective of sign, we find that the amplitudes of the current microjumps are, in most cases, considerably larger than those reported by DA95, suggesting a better resolution of the microevents. This is particularly true for $\\Delta \\dot{\\nu}$, where our measurements could be up to a factor of 4 larger than those observed by DA95. As a consequence, most of the candidate microglitch events found to be insignificant in DA95 were observed to be significant in current analysis.  The unprecedented improvements in both the number and amplitudes of pulsar microglitches could be attributed to some unique features of the current timing data: long time coverage ($\\sim$ 16 yrs) and shorter sampling intervals. Shorter data sampling intervals would naturally lead to better time resolutions of events. Expectedly, this would reduce the uncertainty in event epoch, resulting in significantly improved estimates of jump sizes. Glitch analyses (e.g.\\ Flanagan 1990) have demonstrated how the resolution of jump parameters strongly depended on the accuracy of the epochs of the events. The rise times of the events ($\\Delta{t}$) for current analysis lie between 0.05 and 40 d, with a median value of $\\sim$ 10 d. This implies that the epochs of the microglitches reported in this paper are improved by a factor of, at least, 4 over those of JPL and Tasmanian data sets, where  $\\Delta{t}$ are in the range of $\\sim$ 6 $-$ 154 d and 13 $-$ 255 d, respectively. This could, almost certainly, lead to better resolution of event sizes. It is noteworthy that our techniques identified most of the small macroglitches reported previously from independent analyses timing data on the pulsars B0740$-$28 and B1240$-$64 (Janssen \\& Stappers 2006; D'Alessandro \\& McCulloch 1997). Again, our results on the jump parameters (epochs, amplitudes and signatures) are remarkably in good agreement with previous results. However, current methods were able to identify almost all the jumps in $\\dot{\\nu}$, which previous analyses could not resolve owing to large uncertainties in event epochs. Hence, the seemingly high sensitivity of the current analysis could be attributed to the dramatic improvement in precision of the events' epochs in HartRAO data set.  A key challenge in microglitch observation analyses has remained the extraordinary difficulty in distinguishing between real and spurious events. Cordes \\& Downs (1985) suggested that microjumps could be identified by demostrating that they are too large to be mere fluctuations produced by accumulation of many, much smaller events of random walk origin. Previous studies (DA95; Cordes et al.\\ 1988; CD85) have relied largely on the signficance tests, using timing noise strength parameters, to isolate significant events. However, Cordes \\& Downs (1985) have highlighted the strong dependence of the current method of testing the significance of microglitch events on the rise time of the event $\\Delta{t}$. For instance, the standard deviations could possibly be over-estimated or under-estimated for very small or very large $\\Delta{t}$, respectively. This could cause leakage of spurious events with large amplitudes or rejection of extremely small-amplitude discrete events. The former scenario might be the case in previous analyses (DA95; CD85), while current data appear to favour the later  scenario. With a median event rise time ($\\Delta{t}$) of $\\sim 10$ d, it is feared that both $\\sigma_{\\Delta{\\nu}}$ and $\\sigma_{\\Delta{\\dot{\\nu}}}$ could have been significantly under-estimated in some cases, causing significant leakage of spurious events into the sample of real discrete events (e.g.\\ Chukwude 2002). However, we are optimistic that the 2-$\\sigma$ standard formal error condition imposed on the size identified microglitches is stringent enough to filter out all or, at least, most spurious events."
    },
    "1005/1005.2584_arXiv.txt": {
        "abstract": "We use different particular classes of axially symmetric Szekeres Swiss-cheese models for the study of the apparent dimming of the supernovae of type Ia. We compare the results with those obtained in the corresponding Lema\\^itre--Tolman Swiss-cheese models. Although the quantitative picture is different the qualitative results are comparable, i.e, one cannot fully explain the dimming of the supernovae using small scale ($\\sim 50$ Mpc) inhomogeneities. To fit successfully the data we need structures of order of 500 Mpc size or larger. However, this result might be an artifact due to the use of axial light rays in axially symmetric models. Anyhow, this work is a first step in trying to use Szekeres Swiss-cheese models in cosmology and it will be followed by the study of more physical models with still less symmetry. ",
        "introduction": "The Universe, as we observe it, is inhomogeneous almost at all scales. However, the standard approach, based on the cosmological principle, assumes that homogeneous models with Robertson-Walker (RW) geometry can be successfully employed in a cosmological framework, i.e., above some $\\sim$100 Mpc scale. However, within such a framework, one needs to assume that the Universe is filled with dark energy in order to achieve concordance with SNe Ia observations. But dark energy has never been directly observed neither in the Universe, nor in laboratories, and it has very exotic properties, namely, it is a kind of gas with negative pressure or a mere cosmological constant with an amplitude too small to account for the vacuum energy in the standard model of particle physics. Therefore, it is of the utmost importance to test whether the observations can be explained with models that do not require dark energy but that instead take inhomogeneities into account.  A number of approaches have been suggested so far. Among them the one patch spherically symmetric Lema\\^itre--Tolman models with a central observer \\cite{celerier00, iguchi02, alnes06, chung06, alnes07, enqvist07, ABNV09, bolejko08a, garciabellido08, garciabellido08b,zibin08,yoo08,enqvist08, BW09,CBKH09}, the thin shell approximations \\cite{tomita00,tomita01a,tomita01b,CS08}, the Stephani models \\cite{dabrowski98,godlowski04,stelmach06,dabrowski07}, and the Szekeres models \\cite{MI1,MI2}. However, all these approaches have their restrictions.    For example, for the Lema\\^itre-Tolman model assuming spatial spherical symmetry around the observer, this symmetry has been turned into an argument stating that the observer's location at or near the symmetry center is too special as regards the requirement of the cosmological principle. Now, we want to stress here that these models should not be considered as assuming we are living at or near the center of any spherically symmetric universe. They should merely be considered as a first step in the process of modeling cosmological inhomogeneities, i.e., as a mere smoothing out of the inhomogeneities on angular scales, analogous to their smoothing out on the totality of space in homogeneous models. Thus, the subsisting inhomogeneities are only radial. The use of such models must be merely regarded as a first approach which will be followed by more precise ways of dealing with the observed inhomogeneities. The Stephani class of solutions are conformally flat \\cite{Krasinski97} and cannot therefore constitute general models for the Universe. The thin shell approximations are based on two RW models and are also spherically symmetric around the observer. The Szekeres models are the most general ones, even if the five functions defining a model depend only on the radial coordinate.   An improvement is to use Swiss-Cheese models. Some models of this class where the holes were Lema\\^itre--Tolman patches have already been proposed in the literature \\cite{BTT07,MKMR07,ABNV09,BN08,BTT08,KMM09}. However, the results suggest that we need at least structures of several hundred Mpc in order to reproduce the supernova data. In Ref. \\cite{CF09} the model proposed contains Schwarzschild regions, each with the mass of a galaxy. The results suggest that such a geometry cannot explain dark energy as an effect of inhomogeneities.   We propose here a generalization of these approaches, where instead of Lema\\^itre--Tolman or Schwarzschild patches, which are all spherically symmetric, we use Szekeres models to represent the holes in the homogeneous cheese. The Szekeres model has several advantages, like anisotropic density distribution which leads to a different evolution pattern than in the Lema\\^itre--Tolman model \\cite{B06,B07}. Thus there is a possibility that the Szekeres Swiss-cheese models can come up with a more realistic picture, where the inhomogeneities are no more spherical and therefore can give a more physical representation of our Universe. However, for calculation simplicity and as a first step, we limit ourselves here to the study of an axially symmetric propagation of light through axially symmetric Szekeres holes.  The structure of the present paper is as follows. Section \\ref{sec3} is devoted to the presentation of the Szekeres models and of the particular subclasses used in this paper. In Sec.~\\ref{sec4}, we give the residual Hubble diagrams for the five axially symmetric quasi-spherical Szekeres Swiss-cheeses considered here and we compare them to the corresponding Lema\\^itre--Tolman models, to the $\\Lambda$CDM model and to the actual data. In Sec.~\\ref{sec5}, we display our conclusions. ",
        "conclusions": "\\label{sec5}  We have used the quasi-spherical axially symmetric subclass of the Szekeres models to construct Swiss-cheese cosmological models which are generalizations of other Swiss-cheeses where the holes are spherically symmetric and that have already been proposed in the literature. This geometry has already proved to be very useful when studying structure formation as they allow for matter anisotropies \\cite{B06,B07}.  However, as regards supernova observations, we do not get qualitatively different pictures from the ones obtained with the corresponding Lema\\^itre--Tolman Swiss-cheese models. The results given in Sec. \\ref{results} show that, as in Lema\\^itre--Tolman models \\cite{BTT07, BTT08, KMM09}, small-scale inhomogeneities ($\\sim 50$ Mpc) do not alter the distance-redshift relation significantly, and thus cannot fully explain away dark energy. To reproduce successfully the supernova dimming, we must consider larger structures of order 500 Mpc, as in Lema\\^itre--Tolman Swiss-cheeses \\cite{ABNV09}. The presence of such large structures, in the framework of Lema\\^itre--Tolman models, can be ruled out by CMB constraints \\cite{WV09}. However, for Szekeres models, the constraints might be different since, as shown in \\cite{KB09}, the CMB fluctuations are smaller than for the Lema\\^itre--Tolman Swiss-cheese models.  The fact that in the studied cases we need large patches in order to explain supernova dimming without dark energy might be due to the axial symmetry of the models. Actually, it has been shown, that structure formation is five times faster in some non axially symmetric Szekeres models than in the corresponding Lema\\^itre--Tolman models \\cite{B06}. Since we know, from \\cite{MKMR07}, that the evolution of the inhomogeneities bends more the photon paths compared to the FLRW case than their mere geometry, it should be interesting to investigate if another geometry of the holes might be able to enhance their evolution and therefore to have a stronger effect on the bending of the photon paths and, hence, on the supernova dimming. This will be the subject of future work."
    },
    "1005/1005.2298_arXiv.txt": {
        "abstract": "\\vspace{10mm} \\begin{center} \\textbf{Abstract} \\end{center} We discuss a new mechanism which can be responsible for the origin of the primordial perturbation in inflationary models, the inhomogeneous DBI reheating scenario. Light DBI fields fluctuate during inflation, and finally create the density perturbations through modulation of the inflation decay rate. In this note, we investigate the curvature perturbation and its non-Gaussianity from this new mechanism. Presenting generalized expressions for them, we show that the curvature perturbation not only depends on the particular process of decay but is also dependent on the sound speed $c_s$ from the DBI action. More interestingly we find that the non-Gaussianity parameter $f_{NL}$ is independent of $c_s$. As an application we exemplify some decay processes which give a viable and detectable non-Gaussianity. Finally we find a possible connection between our model and the DBI-Curvaton mechanism. ",
        "introduction": "To account for the density perturbation which seeds the structure of the observed universe, the inflationary paradigm\\cite{KT} is a promising candidate. In this picture, the universe went through an accelerated expansion in the very early period. This scenario predicts that the inflaton field $\\phi$ rolls downs its potential with quantum fluctuations superimposed which lead to density perturbations. However it is important to investigate alternatives to this simple scenario.  Recently more and more evidence from observations of cosmic microwave background (CMB) anisotropies favors that the primordial density fluctuations (PDF) are almost Gaussian, Scale invariant, and adiabatic \\cite{Spergel,Komatsu:2010fb}. And, according to many works \\cite{Lyth:2001nq,Moroi:2001ct,Dimopoulos:2002kt,Lyth:2002my,Lyth:2005ab,Lyth:2005fi,Dvali:2003em,Kofman:2003nx,Dvali:2003ar,Zaldarriaga:2003my,Matarrese:2003tk,Suyama:2007bg,Battefeld:2007st,Ichikawa:2008ne}, light scalar fields merit attention to investigate their role as candidates for explaining the origin of the PDF. Such light scalar fields generally exist in extensions of the standard model of particle physics, which motivates some alternatives. One is the curvaton scenario \\cite{Lyth:2001nq,Moroi:2001ct,Dimopoulos:2002kt,Lyth:2002my,Lyth:2005ab,Lyth:2005fi}, in which the final curvature perturbations are produced from an initial isocurvature perturbation associated with the quantum fluctuations of a light scalar field other than the inflaton, the curvaton, whose energy density is negligible during inflation. The curvaton isocurvature perturbations are transformed into adiabatic ones when the curvaton decays into radiation well after the end of inflation. Another is the inhomogeneous/modulated reheating scenario \\cite{Dvali:2003ar,Dvali:2003em,Zaldarriaga:2003my,Kofman:2003nx,Matarrese:2003tk,Vernizzi:2003vs}. This supposes that the decay rate $\\Gamma$ of the inflaton varied in space due to a dependence on a light field, and density perturbations would be generated during reheating independently of those generated by the standard inflationary mechanism. In these two scenarios the light scalar fields, not the inflaton, are responsible for the primarily density perturbations. The inflaton just serves to drive and end inflation, and under this assumption the constraints on inflation are considerably lessened. By introducing these light scalar fields one also finds that primordial non-Gaussianity can be very large compared to the single inflaton models. In curvaton models the non-linear parameter $f_{NL} \\sim 5/4r$ where $r$ is a small coefficient \\cite{Lyth:2002my,Lyth:2005ab,Lyth:2005fi}; and inhomogeneous/modulated reheating models suggest $f_{NL} \\sim O(1)$ \\cite{Dvali:2003ar,Zaldarriaga:2003my,Matarrese:2003tk} or larger given particular decay processes, see Refs \\cite{Suyama:2007bg,Battefeld:2007st,Ichikawa:2008ne,Kohri:2009ac}.  Cosmological modellers have paid attention recently to the DBI field as an alternative to a canonical field. Some earlier works, Refs.\\cite{Silverstein:2003hf,Alishahiha:2004eh,Chen:2006nt,Easson:2007dh,Easson:2007fz,Gmeiner:2007uw,Arroja:2008yy}, have discussed cosmology using a DBI field, and with this interesting source some results are different from those obtained in Refs. \\cite{Huang:2007hh,Langlois:2008mn,Langlois:2008wt,Langlois:2008qf,Guo:2008sz}. So far only the ordinary light scalar fields were discussed in inhomogeneous reheating, therefore in this note, we focus on the inhomogeneous(modulated) reheating models specifically involving DBI fields. As a toy model to reheat the universe, we assume that the DBI fields can dominate the decay rate. With the variation of the decay rate, the fluctuations can be generated after inflation. ",
        "conclusions": "\\subsection{Observable Parameters} As mentioned above, we have found some results which differ from previous works. From Eq.(\\ref{dbi-inho}), we can see that the curvature perturbation becomes ${\\cal{P}}^{1/2}_{\\zeta} \\sim {\\frac{\\sqrt{c_s}}{12\\pi}}\\frac{H_*}{M_*}$; here $\\rm M_*$ is the mass scale during reheating. In the inflationary background, gravitational waves can be predicted as the tensor fluctuation ${\\cal{P}}_{T}^{1/2} \\sim {\\frac{1}{\\pi} \\frac{H_*}{M_{\\rm pl}}}$.  \\textbf{(1)} At first glance, due to the existence of the sound speed the primordial curvature perturbation, $\\zeta$ in Eq.(\\ref{dbi-inho}), can be suppressed by the term $\\sqrt{c_s}$, but we also note that the detailed decay process plays a significant role in determining the final value of the curvature perturbation.  \\textbf{(2)} Considering the ratio $r$ of tensor to scalar, here in our model it is given by \\be r = \\frac{{\\cal P}_{T}}{{\\cal P}_{\\zeta} }\\sim 10^2 \\frac{1}{c_s} \\frac{f^2}{M_{pl}^2 f'^2} \\ee Like any linear approximation, in expanding we assume that the decay process has the form of $f(\\sigma) \\propto (1 + \\mathrm{\\vartheta \\sigma/M} + ... )$ (similarly see Refs.\\cite{Dvali:2003em,Ichikawa:2008ne}), where $\\vartheta$ is small, and $\\rm M$ (by ignoring subindex $*$) is a dimensional parameter representing the mass scale for reheating after inflation ending, and $m_\\sigma < M < M_{pl}$. If we use $r < 0.36 (95\\% \\rm{CL})$ from 7-year WMAP limit in the Ref.\\cite{Komatsu:2010fb}, then we should impose $f'/f \\sim \\vartheta/M_*$, so that $M_* \\sim \\vartheta \\sqrt{c_s} M_{\\rm pl} \\times 10^{-3}$. This means that for very small $c_s$, we can get the $m_\\sigma \\ll M_{\\rm pl}$ which is assumed in the beginning of our consideration.  \\textbf{(3)} Contrary to the standard inhomogeneous reheating scenario of Ref.\\cite{Dvali:2003em}, in our model the created fluctuations are of order $\\sqrt{c_s} H_* \\frac{f'}{f} \\sim \\vartheta \\sqrt{c_s} \\frac{H_*}{M_*}$. The relation between the slope of the power spectrum of density perturbations and the inflaton potential in various scenarios is, \\ba n-1 &=& \\frac{d \\ln{ \\Big((f'/f)^2c_sH^2}\\Big) }{d \\ln{a}}\\quad \\quad \\quad \\mbox{in our model} \\nonumber \\\\ n-1 &=& \\frac{d \\ln{H^2}}{d \\ln {a}}\\quad \\quad \\quad \\quad \\quad \\quad \\quad \\quad \\mbox{standard inhomogeneous reheating} \\nonumber \\\\ n-1 &=& \\frac{d \\ln{H^2/\\epsilon}}{d \\ln{a}}\\quad \\quad \\quad \\quad \\quad \\quad \\quad  \\mbox{standard inflation scenario} \\ea Our result is different from both Ref.\\cite{Dvali:2003em} and the standard inflationary scenario.  Discussing the spectral index explores to what degree the decay process and the warp space contribute to it. With $c_s = \\sqrt{1-2Xh(\\sigma)}$ and its derivative with respective to DBI field $\\rm \\sigma$, i.e. $c_s^\\prime = \\frac{(c_s^2 - 1)}{2c_s} \\frac{h'(\\sigma)}{h(\\sigma)}$, then we obtain \\ba \\label{h-van} n-1 = \\frac{d \\ln{H_*^2}}{d\\ln a} + \\left[ \\frac{c_s^2 - 1}{2c_s^2} \\frac{h'(\\sigma)}{h(\\sigma)} + 2 \\frac{f'(\\sigma)}{f(\\sigma)} \\left(\\frac{f(\\sigma)f''(\\sigma)}{f'^2(\\sigma)} - 1\\right)  \\right] \\frac{d\\sigma}{d\\ln{a}} \\ea where parameter $h(\\sigma)$ represents the warped throat and $f(\\sigma)$ indicates the decay rate. For simplicity for estimating the contribution to power spectra, we set all $h'(\\sigma)/h(\\sigma)$ and $f'(\\sigma)/f(\\sigma)$ as proportional to $1/\\sigma$ for easy estimation. By doing this, we can combine the second term as \\ba n-1 &\\sim& \\frac{d \\ln{H_*^2}}{d\\ln a} + \\left[ \\frac{c_s^2 - 1}{2c_s^2} + 2 \\left(\\frac{f(\\sigma)f''(\\sigma)}{f'^2(\\sigma)} - 1\\right)  \\right] \\frac{d\\ln{\\sigma}}{d\\ln{a}} \\nonumber\\\\ &\\sim& \\frac{d \\ln{H_*^2}}{d\\ln a} + \\left( \\frac{c_s^2 - 1}{2c_s^2} + \\frac{2}{5} f_{NL}  \\right) \\frac{d\\ln{\\sigma}}{d\\ln{a}} \\ea  \\textbf{(A):} if $c_s \\simeq 1$, the decay rate surpasses the contribution from the warped throat. From Eq.(\\ref{h-van}), the effect of the warp throat is reduced. However, in view of the Lagrangian, the system is just reduced to a canonical one, and can be rewritten as a usual dynamic system with $\\mathcal{L} = \\mathcal{L} {(X,\\sigma)} = {-\\frac{1}{2}{\\partial_{\\mu}{\\sigma}\\partial^{\\mu}{\\sigma}} - V(\\sigma)}$.  \\textbf{(B):} if $c_s \\ll 1$, it is hard to say which contribution affects the whole term more. Because of the warped factor $h$, the sound speed $c_s$ and the unknown decay process $f(\\sigma)$, the model possesses a lot of freedom, but reheating from a warped throat is worthy of investigation in future work.  \\subsection{Some Specific Examples} Interestingly, from Eq.(\\ref{in-fnl}), the non-Gaussianity is independent of the sound speed in our model even though the curvature perturbation depends on it. This distinguishes from the previous works in DBI field scenario \\cite{Huang:2007hh,Langlois:2008mn,Langlois:2008qf,Li:2008fm} where the primordial fluctuations and non-Gaussianities are both enhanced by a low sound speed. Here the non-Gaussianities only depend on the specific decay process.  We wish to know whether a particular decay process exists to generate a large curvature perturbation and also non-Gaussianities. Let's consider some specific toy models able to produce detectable density perturbations and large non-Gaussianity.  \\begin{figure}[ht] \\setlength{\\abovecaptionskip}{0pt} % \\begin{minipage}[b]{0.48\\linewidth} % \\includegraphics[width=2.8in]{toy-log-pert3.eps} \\hspace{0.4cm} \\label{fig:zeta} \\end{minipage} \\begin{minipage}[b]{0.48\\linewidth} % \\includegraphics[width=2.8in]{toy-log-fnl3.eps}% \\label{fig:fnl} \\end{minipage} \\caption{(Left) Plot of the ratio $\\mathcal \\Re = P_\\zeta / P_*$, where ${P_*}$ indicates the scale-invariant power spectrum $P_* = \\left(\\frac{H_*}{2\\pi}\\right)^2$. (Right) non-Gaussianity of perturbation. The plots show that the amplitude of $f_{NL}$ is of O(1) , even though $\\sigma \\ll { M}$. There, the mass scale for $M$ is lower than inflaton scale ${H_*}$.}\\label{fig:Curvature-fnl} \\end{figure} \\textbf{(1)} We assume that the decay process is described by $f(\\sigma) \\sim \\Big(1 + 6 \\log \\left(1 + \\sigma/{\\rm M} \\right) \\Big)^{-1}$. Here $\\rm M$ is a mass scale during reheating, which is less than $\\rm H_*$.  In this case with Fig.1, we find $\\zeta \\sim \\frac{\\sqrt{c_s}}{(1 + 6\\log (1+\\sigma/{\\rm M}))(1+\\sigma/ {\\rm M})}$ and $f_{NL} = \\frac{35}{6} + 5\\log(1+\\sigma/ {\\rm M})$. From the plot on the left hand side, when the speed of DBI field $\\sigma$ reaches its relativistic limit, which means the sound speed $c_s$ approaches zero, we get the curvature perturbation converted from the fluctuation of $\\sigma$ to be zero. If the sound speed is approaching $1$, meaning corresponding to a canonical system, the curvature perturbation can be of order $P_*$. However the ratio $\\mathcal \\Re$ is less than $1$, from the left plot, generally the curvature perturbation of the DBI field through decaying process is suppressed by $c_s$, even if the DBI field has a large mass scale.  We can see the non-Gaussianity from the plot on the right hand side. It is independent of the sound speed $c_s$. It is possible to obtain a large non-Gaussianity, being of O(10), if we place suitable coefficients in the decay rate.  \\textbf{(2)} If we take for example the form of $\\left(1 - \\lambda  \\sigma / M \\right)^{-1/n}$, we could get $f_{NL} \\sim O(1 \\sim 10)$ or more. Here $\\lambda$ is just a sign which is used to determine the curvature perturbation together with the power $n$, while $M$ is a mass scale during reheating. \\begin{figure}[ht] \\begin{minipage}[b]{0.48\\linewidth} % \\includegraphics[width=2.8in]{toy-lin-pertn3.eps}% \\label{fig:n} \\end{minipage} \\begin{minipage}[b]{0.48\\linewidth} % \\hspace{0.4cm} \\includegraphics[width=2.8in]{toy-lin-pert-n3.eps}% \\label{fig:-n} \\end{minipage} \\caption{ Different shapes of the ratio $\\mathcal \\Re = \\frac{P_\\zeta}{P_*}$ with $n>0$ (Left) and $n<0$ (Right). Here we place $n=1$. When $\\left|n\\right|$ is larger than one, the perturbation will be more flat.} \\label{fig:Curvature-fnl} \\end{figure}  In Fig. 2, we notice the curvature perturbations have different shapes in case with different $n$ and the sound speed $c_s$ overall. The figures show that the perturbations are smoothed in different parametric regions. The non-Gaussianity $f_{NL}$ of these types of perturbation has the same shape, with $f_{NL} \\equiv 5n$ regardless of the sign of $n$, even given the different $c_s$.  If $c_s \\sim 0$, the curvature perturbation approaches zero too. When $c_s$ gets larger, we can find the possibility of generating a detectable perturbation when $n>0$ (in the left figure). In contrary, when $n<0$ (in the right figure), it might not easy to detect the curvature perturbation which is converted from the DBI field during the decay period."
    },
    "1005/1005.0254_arXiv.txt": {
        "abstract": "{ We examine the nonlinear magnetofrictional extrapolation scheme using the solar active region model by Titov and D\\'emoulin as test field. This model consists of an arched, line-tied current channel held in force-free equilibrium by the potential field of a bipolar flux distribution in the bottom boundary. A modified version, having a parabolic current density profile, is employed here. We find that the equilibrium is reconstructed with very high accuracy in a representative range of parameter space, using only the vector field in the bottom boundary as input. Structural features formed in the interface between the flux rope and the surrounding arcade---``hyperbolic flux tube'' and ``bald patch separatrix surface''---are reliably reproduced, as are the flux rope twist and the energy and helicity of the configuration. This demonstrates that force-free fields containing these basic structural elements of solar active regions can be obtained by extrapolation. The influence of the chosen initial condition on the accuracy of reconstruction is also addressed, confirming that the initial field that best matches the external potential field of the model quite naturally leads to the best reconstruction. Extrapolating the magnetogram of a Titov-D\\'emoulin equilibrium in the unstable range of parameter space yields a sequence of two opposing evolutionary phases which clearly indicate the unstable nature of the configuration: a partial buildup of the flux rope with rising free energy is followed by destruction of the rope, losing most of the free energy. } ",
        "introduction": "\\label{s:intro}  The computation of the coronal magnetic field from boundary data is the only available technique to obtain fully three-dimensional information, i.e., the complete field vector in the whole coronal volume bounded below by the magnetogram. Deductions of the field from spectral lines formed in the coronal plasma, from radio maps across a range of microwave frequencies, or from observations of specific structures such as prominences, are restricted to two-dimensional projections or to cuts through the volume, or do not yield the full vector. The three-dimensional information is required to advance a wide array of issues in coronal physics, for example the structure of prominences, the onset of eruptions, and how the coronal field is connected to the Sun's interior and the solar wind.  In equilibrium, the coronal magnetic field is nearly force-free in the major part of active regions, except their upper periphery \\cite[see, \\eg][]{1995ApJ...439..474M, 2002ApJ...568..422M, 2001SoPh..203...71G}. Hence, to a very good approximation it can be described by \\begin{equation} \\Nabla \\times \\vB = \\alpha \\vB \\qquad \\mbox{with} \\qquad \\Nabla \\cdot \\vB=0. \\label{eq:fff} \\end{equation} Here $\\alpha(\\vx)$ is a scalar function that is constant along each individual field line (which follows directly from Eqs.~[\\ref{eq:fff}]) but in general has different values on different field lines.  The computation of force-free coronal fields can be formulated as a boundary-value problem for Eqs.~(\\ref{eq:fff}). Such extrapolation can be performed at different levels of approximation for the field. If $\\alpha$ is assumed to be constant, then the first of Eqs.~(\\ref{eq:fff}) is linear, and only the knowledge of the normal field component is needed as boundary condition for the extrapolation. The solution of the linear problem was given in closed form in \\cite{1972SoPh...25..127N, 1977ApJ...212..873C, 1978SoPh...58..215S}. Linear methods have several intrinsic limitations, the most severe one being that they cannot match most observed magnetograms closely because $\\alpha$ is often found to vary strongly across active regions \\citep[\\eg][]{2002A&A...392.1119R}. The particular case of $\\alpha=0$ yields a potential field, which is inadequate for many purposes, since it does not contain any free energy to power coronal activity. Therefore, it is necessary to proceed beyond the linear approximation and to solve the nonlinear extrapolation problem, permitting $\\alpha(\\vx)$ to vary across the field \\citep{1981SoPh...69..343S,1997SoPh..174..191M,1999A&A...350.1051A}. Due to the nonlinearity, such extrapolation requires a numerical approach. Various numerical schemes and implementations have been developed to construct a solution; for a recent review see \\citet{2006SoPh..235..161S}.  A partially alternative strategy, the flux rope insertion method \\citep{2004ApJ...612..519V, ballegooijen07}, permits to model structures that are supposed to contain a flux rope, particularly filament channels. Here an extrapolated potential field is replaced by a flux rope in part of the volume, and the resulting configuration is numerically relaxed to find an equilibrium. The parameters of the inserted rope are varied until the field lines of the resulting equilibrium match observed features, e.g., threads in a filament or overlying loops.  A convenient tool for both nonlinear force-free extrapolation and the flux rope insertion method is the so-called magnetofrictional relaxation of the field to force-free numerical equilibria \\citep{1981JCoPh..41...68C, 1986ApJ...309..383Y, 1986ApJ...311..451C, 1996ApJ...473.1095R}. In a previous paper \\citep{2005A&A...433..335V}, we described our implementation of magnetofrictional nonlinear extrapolation and its application to reconstruct a relatively simple force-free magnetic field, which contains an approximately current-neutralized flux rope rooted in the flux concentrations of a simple bipolar magnetogram \\citep{2003A&A...406.1043T}. That paper included a thorough study of the dependence of the reconstruction quality on the numerical resolution. In \\cite{2007SoPh..245..263V} an improved implementation of the magnetofrictional method allowed us to significantly increase the accuracy in the reconstruction of the well known force-free equilibria by \\cite{1990ApJ...352..343L} in comparison to similar extrapolations summarized in \\cite{2006SoPh..235..161S}. Applications of this code to measured magnetograms are included in \\cite{2008SoPh..247..269M}, \\cite{2008ApJ...675.1637S}, and \\cite{2009ApJ...696.1780D}.  Here we extend the testing of our magnetofrictional code in the reconstruction of solar-relevant solutions of Eqs.~(\\ref{eq:fff}) by applying it to the flux rope equilibrium constructed by \\citeauthor{1999A&A...351..707T} (\\citeyear{1999A&A...351..707T}, hereafter TD). This configuration describes structures that emerge from the solar interior already twisted and carry a net current \\cite[e.g.,][]{1995ApJ...446..877L, 1996ApJ...462..547L, 2000ApJ...532..616W}, and it is in this sense complementary to the one in \\cite{2005A&A...433..335V} which represents the case of an originally potential field twisted solely by photospheric displacements. The relevance of the TD equilibrium as a model for solar active regions and eruptive configurations was demonstrated by many investigations \\citep[\\eg][]{2000SPD....31.0236Z,2004A&A...413L..23K,2005ApJ...630L..97T, 2008ApJ...674..586S, 2008A&A...481L..65M}. In particular, it has been shown that the TD flux rope can be subject to ideal MHD instabilities \\citep{2003ApJ...588L..45R,2004A&A...413L..27T,2007AN....328..743T,2007ApJ...670.1453I}. Depending on its parameters, the equilibrium can include a \\emph{hyperbolic flux tube (HFT)} or \\emph{bald patches (BPs)} (see Sect.~\\ref{s:tdeq} for detail). Such \\emph{structural features} are thought to play an important role in the initiation and evolution of coronal mass ejections (CMEs) and their associated flares. Moreover, the occurrence of BPs changes the topology of the field.  Our goal in this paper is to provide a detailed analysis of the reconstruction capabilities of the magnetofrictional code for this test field. We address the flux rope morphology, including the twist and the structural features HFT and BP, the energy and helicity contents, and the influence of the initial condition used by the extrapolation code. We apply the code also to an unstable case, expecting this to be useful in judging extrapolations of magnetograms that were taken during the initial phase of an eruption or immediately prior to its onset. Finally, we will compare our results to the previous reconstruction of the TD field by \\cite{2006A&A...453..737W}, which employed a different extrapolation scheme.  Using a force-free test field implies that the magnetogram is compatible with the force-free hypothesis of the extrapolation method. Therefore, in the following we do not need to consider deviations from force-freeness in the photosphere, the ambiguity of the transverse field direction, noise, and insufficient resolution, all of which are associated with measured magnetograms. Attempts to deal with some of these problems are discussed, e.g., in \\cite{1997SoPh..174..191M}, \\cite{1997SoPh..174..129A}, \\cite{2006SoPh..233..215W}, \\cite{2007A&A...476..349F}, \\cite{2008SoPh..247..269M}, \\cite{2008SoPh..247..249W}, \\cite{2008ApJ...675.1637S}, and \\cite{2009ApJ...696.1780D}.  The paper is organized as follows. The TD equilibria used below are described in Sect.~\\ref{s:tdeq}, the extrapolation method is briefly outlined in Sect.~\\ref{s:method}, and the results are presented and discussed in Sect.~\\ref{s:results}. Our conclusions are summarized in Sect.~\\ref{s:conclusions}. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{alpha_modb.eps}} \\caption{``Line-of-sight'' magnetogram $B_z(x,y,0)$ with superimposed arrows representing the transverse field \\emph{(top)}, and force-free parameter $\\alpha(x,y,0)$ \\emph{(bottom)}, for the reference fields of the cases High\\_HFT \\emph{(left)} and BP \\emph{(right)}. \\label{f:alpha_modb}} \\end{figure} ",
        "conclusions": "\\label{s:conclusions}  This investigation demonstrates that a force-free model of solar active regions, which includes relevant structural (i.e., topological or geometrical) features and is free of noise, can be reconstructed to a very high accuracy based only on its vector magnetogram. The reconstructed configuration consists of a flux rope with a non-neutralized current channel in its core. A bald patch separatrix surface or two quasi-separatrix layers combined into a hyperbolic flux tube form the interface between the flux rope and the ambient potential field. These elements are regarded to be generic for many active regions, so that the model realistically captures their structural properties, in particular for essentially bipolar flux distributions.  Previous reconstructions of nonlinear force-free test fields considered the simpler configurations of an approximately current-neutralized flux rope \\citep{2005A&A...433..335V} and of a sheared arcade \\citep{2006SoPh..235..161S, 2007SoPh..245..263V}. A force-free model of a filament channel containing a hollow current channel could also be reconstructed, albeit at significantly lower accuracy, which was likely due to the combination with a noisy line-of-sight magnetogram in the exterior of the channel \\citep{2008SoPh..247..269M}. Thus, the reconstruction capability has now been demonstrated, mostly at excellent accuracy, for all basic types of force-free equilibria currently considered in coronal physics, with the one treated here providing the highest degree of structural complexity. This confirms and substantiates previous conclusions \\citep{2008SoPh..247..269M, 2009ApJ...696.1780D} that insufficient accuracy of nonlinear force-free extrapolation originates from the violation of the force-free assumption by the input vector magnetogram, not from insufficient intrinsic accuracy of the extrapolation codes.  The accuracy achieved in the reconstruction of stable modified TD equilibria clearly permits to discriminate between configurations containing an HFT, a BP, or none of these features. The energy, flux rope twist and position, and helicity are obtained at relative error levels of $\\la$\\,1, 3, 4, and $\\sim\\!5\\%$, respectively, even for a high twist near the threshold of the helical kink instability, slightly exceeding one full turn.  These quantities are found to depend only weakly on the initial field chosen, except for the apex height of the flux rope which shows a displacement from the true value of up to $\\approx20\\%$ (if a linear force-free field is used). Starting the extrapolation from the conventional potential field, which assumes vanishing flux outside the magnetogram, yields the most accurate results for the considered equilibria. Other potential fields result in somewhat lower accuracy, and a linear force-free field scores worse but still produces a rather acceptable reconstruction.  The high accuracy is reached using vector magnetograms that have a realistic low-flux margin of width $\\sim1/4$ the size of the strong-flux area and are exactly flux balanced. The effects that result from violating the latter condition require further study.  Attempting to extrapolate the magnetogram of an unstable modified TD equilibrium, the magnetofrictional relaxation produces an evolution in two phases of opposing tendency which clearly reflect the unstable nature of the equilibrium. A phase of rising free energy yielding a partial buildup of the flux rope is superseded by a decrease of the energy back to nearly the initial (potential-field) value, destroying the flux rope."
    },
    "1005/1005.5701_arXiv.txt": {
        "abstract": "We use a recent sample of 49 galaxies to show that there is a proportionality relation between the black hole mass $M_{\\rm BH}$ and the quantity $\\mu\\equiv M_{\\rm G}\\sigma/c$, where $M_{\\rm G}$ is mass of the spheroidal stellar component and $\\sigma$ is the stellar velocity dispersion. $\\mu$ is called the {\\it momentum parameter} and the ratio is $M_{\\rm BH} / \\mu \\approx 3.3$. This result is applied to the penetrating-jet feedback model which argues that the correlation that holds is with a momentum-like parameter, although this feedback mechanism is based on energy balance. ",
        "introduction": " ",
        "conclusions": "We examined the correlation between $M_{\\rm BH}$, and the momentum parameter, $\\mu\\equiv M_{\\rm G} \\sigma/c$. The motivation for this study is the penetrating-jet feedback mechanism that predicts such a correlation (Section \\ref{sec:model}; \\citealt{Soker2009}), although the SMBH determines the correlation by depositing energy into the ISM rather than momentum. Using the sample of \\citet{Gultekin_etal2009} we examined the correlation with $\\mu$ in Section \\ref{sec:correlation}. We applied our results to the penetrating-jet feedback mechanism in Section \\ref{sec:feedback}. The main results of our study and the implied insights can be summarized as follows: \\begin{enumerate} \\item Despite the large statistical uncertainties, recent data suggest that the masses of SMBHs are indeed correlated with the momentum parameter of the bulge. Moreover, we find that the relation is a linear one. \\item The penetrating-jet feedback mechanism \\citep{Soker2009} is compatible with such a relation. \\item In the jet-penetrating feedback mechanism, $\\eta_p$ is a fundamental parameter that must be obtained from observations. It has the same role as $\\eta$ in equation (\\ref{eq:eta}), the relation for the total energy released by the accreting BH, but for the momenta of the two jets. We find $\\eta_p \\sim 0.04$, which is of course an average value over many systems, and average over time in each system (during the phase of SMBH growth). \\item This allows us to scale equation (\\ref{eq:cor1}) by \\begin{eqnarray} \\frac {M_{\\rm G}} {M_{\\rm BH}}  = 480 \\left( \\frac{\\sigma}{200 \\kms} \\right)^{-1} \\left( \\frac{\\eta_p}{0.04} \\right) \\label{eq:corf} \\end{eqnarray} \\item We further analysed the implication of $\\eta_p \\sim 0.04$, and derived an expression for the fraction of the SMBH power that is carried by the jets (equation \\ref{eq:f}). While relativistic jets carry $\\sim 38$ per cent of the released power, non relativistic jets carry a small fraction of the power (the rest is in radiation). Although the latter only carry a small fraction of the total power, they are still able to maintain the feedback process. In this model, there is no need for ``fine tuning'' in using only a small fraction of the jets' energy, but instead a substantial fraction of the energy is used. \\item We cannot determine with current data which of the correlations of $M_{\\rm BH}$ (with $\\mu$, $\\sigma$, $M_{\\rm G}$, or $M_{\\rm G}\\sigma^2$) is better. \\end{enumerate}"
    },
    "1005/1005.1314_arXiv.txt": {
        "abstract": "We investigate the long-term dynamical evolution of star clusters in a steady tidal field produced by its parent galaxy. In this paper, we focus on the influence of mass profile of the parent galaxy. The previous studies were done with the simplification where the parent galaxy was expressed by point mass. We express different mass profiles of the parent galaxy by the tidal fields in which the ratios of the epicyclic frequency $\\kappa$ to the angular velocity $\\omega$ are different. We compare the mass-loss timescale of star clusters whose tidal radii are identical but in parent galaxies with different mass profile, by means of orbits calculations in fixed cluster potential and N-body simulations. In this situation, a cluster rotates around the parent galaxy more rapidly as the parent galaxy has shallower mass profile. We found that the mass-loss timescale increase $20\\%$ and $50\\%$ for the cases that the mass density profile of the parent galaxies are proportional to $R^{-2}$ and $R^{-1.5}$ where $R$ is the distance from the galaxy center, compared to the point-mass case, in moderately strong tidal field. Counterintuitively, a cluster which rotates around the parent galaxy more rapidly has a longer lifetime. The increase of lifetime is due to the fact that the fraction occupied by regular-like orbit increases in shallower profile. Finally, we derive an evaluation formula for the mass-loss timescale of clusters. Our formula can explain a property of the population of the observed galactic globular clusters that their half-mass radii become smaller as their distances from the galactic center become smaller. ",
        "introduction": "In a star cluster, stars exchange their energy with the other stars through their encounters, and the whole system approaches to thermodynamically relaxed state. This process is called two-body relaxation. In this process, some stars have high energy enough to escape from the cluster and the system loses its mass. Therefore, the mass-loss timescale of a cluster is thought to be proportional to its two-body relaxation time (\\cite{Ambart}; \\cite{Spit1940}).  However, this is not true for a cluster in an external tidal field, which was first shown by means of $N$-body simulations in the Collaborative Experiment (\\cite{Heggie1998}). \\citet{Fuku2000} (hereafter, FH) showed that stars which have enough energy to escape from the cluster (hereafter, ``potential escapers'') can remain inside the cluster before finding exits on timescale comparable to the two-body relaxation time of the cluster, and that the timescale (hereafter, ``escape time delay'') complicates the scaling of mass-loss timescale with two-body relaxation time. \\citet{Baum2001} and \\citet{Baum2003} performed $N$-body simulations in the external tidal field, and found the mass-loss timescale are proportional to the half-mass relaxation time to the power of 3/4. \\citet{Tani2005} (hereafter, paper I) also performed $N$-body simulations, and found that the power depends on the strength of the tidal field and the power 3/4 shown by \\citet{Baum2001} can be observed in the case of moderately strong tide.  These findings mean that for treatment of mass loss of clusters in the external tidal field, it is desirable for orbits of stars to be fully resolved, and $N$-body simulation is more appropriate than orbit-averaged method, such as Fokker-Planck method.  This is a succeeding paper of Paper I. In this paper, we investigate the effect of the mass profile of the parent galaxy which produces the tidal field. In most of previous simulations of star clusters with the external tidal field (\\cite{Giersz94}; \\cite{McMillan94}; \\cite{Fukushige95}; \\cite{Baum2001}; Paper I; \\cite{Trenti07}), the parent galaxy are simplified as point mass.  If we consider the mass profile of the parent galaxy, the equations of motion of star clusters moving in the circular orbit around the galaxy are expressed, with tidal approximation, as \\begin{eqnarray} \\displaystyle \\frac{d^2{\\bf r}_{i}}{dt^2} = - 2 \\left( \\begin{array}{c} 0 \\\\ 0 \\\\ \\omega \\\\ \\end{array} \\right) \\times \\frac{d{\\bf r}_{i}}{dt} - \\left( \\begin{array}{c} (\\kappa^2 - 4 \\omega^2) x_{i} \\\\ 0 \\\\ \\nu^2 z_{i} \\\\ \\end{array} \\right)  - \\nabla \\Phi_{{\\rm c},i},\\label{eq:motion2} \\end{eqnarray} where $\\omega$ is a angular velocity, $\\kappa$ is a epicyclic frequency, $\\nu$ is a vertical frequency, The first term on the right-hand side in equation (\\ref{eq:motion2}) is the Coriolis force, the second term is a combination of centrifugal and tidal forces, and the third term is gravitational acceleration from other particles in the cluster. For the point-mass galaxy, $\\kappa=\\omega$. Also for the galaxy with flat rotation ($\\rho \\propto r^{-2}$), $\\kappa = 2\\omega$. In Table \\ref{tab:rho}, the approximated relations between mass profile and $\\kappa/\\omega$ are summarized. The exact description will be given in section 2.  \\begin{table} \\begin{center} \\caption{Relation among mass density profiles of galaxies, the angular velocities of the clusters, and the epicyclic frequency.} \\begin{tabular}{lll} \\hline \\hline $\\rho\\propto R^{-\\infty}$(point mass) & $\\omega \\propto R^{-3/2}$ & $\\kappa^{2} = \\omega^{2}$ \\\\ $\\rho\\propto R^{-2}$                  & $\\omega \\propto R^{-1}$   & $\\kappa^{2} = 2\\omega^{2}$ \\\\ $\\rho\\propto R^{-1.5}$                & $\\omega \\propto R^{-3/4}$ & $\\kappa^{2} = (5/2)\\omega^{2}$ \\\\ $\\rho\\propto R^{-1}$                  & $\\omega \\propto R^{-1/2}$ & $\\kappa^{2} = 3\\omega^{2}$ \\\\ $\\rho\\propto R^{0}$                   & $\\omega \\propto R^{0}$    & $\\kappa^{2} = 4\\omega^{2}$ \\\\ \\hline \\hline \\end{tabular} \\label{tab:rho} \\end{center} \\end{table}  In this paper, we compare dynamical evolution of the star cluster in the tidal fields of patent galaxies with four different $\\kappa^2/\\omega^2 (=1, 2, 2.5, 3)$, i.e. different mass profile. Throughout the paper, we set the initial tidal radii identically among clusters with different $\\kappa/\\omega$. The tidal radius and that at the initial time are, respectively, expressed as \\begin{equation} r_{\\rm t} = \\left( \\frac{GM}{4\\omega^2 - \\kappa^2} \\right)^{1/3}, \\label{eq:rt} \\end{equation} and \\begin{equation} r_{\\rm t,i}=(GM_{\\rm i}/(4\\omega^2-\\kappa^2))^{1/3}, \\label{eq:rti} \\end{equation} where $G$ is the gravitational constant, $M$ and $M_{\\rm i}$ are the cluster mass and that at the initial time. In the condition, the Coriolis force is relatively stronger for larger $\\kappa^2/\\omega^2$, i.e. shallower mass profile. Although several studies set parent galaxy to have mass density profile other than point mass (\\cite{por2001a}; \\cite{por2002}; \\cite{por2003}; \\cite{Moore1996}; \\cite{Baum2003}; \\cite{por2001b}; \\cite{Yim2002}; \\cite{Dehnen2004}; \\cite{Capuzzo2005}; \\cite{Combes1999}), they aim to study disruption of star clusters taking place in relatively short timescale.  On contrary to them, we investigate systematically the effects of the difference of mass density profile of parent galaxy on long-term mass-loss timescale of clusters .  The plan of this paper is as follows.  In section 2, we give an exact description on the equation of motion in the tidal field.  In section 3, we first perform orbital calculation in the fixed potential of star clusters and we compare the escape time delays of the potential escapers, which complicates the scaling with two-body relaxation timescale.  In section 4, we perform $N$-body simulation of clusters in the tidal field in order to study the influence of mass density profiles of the parent galaxy on the mass-loss timescale of the star clusters. In section 5, we apply our simulation results for galactic globular clusters. Finally we summarize our paper in section 6. ",
        "conclusions": "We investigate mass loss of star clusters in a tidal field of parent galaxy with different mass density profiles. We find that orbit families are different among parent galaxy with different mass density profiles. A fraction of regular orbits increases in the orbit family as the mass density profile of the parent galaxy is shallower. The tendency is strong in nearly planar orbits. The difference of the orbit families affects the distribution of escape time delay in each parent galaxy model. The escape time delay becomes larger as the mass density profile of the parent galaxy is shallower.  We conform that the difference of escape time delay affects mass-loss timescale of clusters with two kinds of size of tidal radii. The mass-loss timescale of clusters becomes larger as the mass density profile of the parent galaxy becomes shallower. In a moderate tidal radius, $r_{\\rm t,i}/r_{\\rm kg}=1.0$, the clusters in parent galaxy with the three shallower mass density profiles ($\\kappa^2/\\omega^2=2,2.5,3$) have larger mass-loss timescale than those in the parent galaxy with the steepest one ($\\kappa^2/\\omega^2=1$) by $20$ \\%, $50$ \\% and a factor of $2.5$, respectively. In the case of larger tidal radius, $r_{\\rm t,i}/r_{\\rm kg}=2.2$, the difference is smaller and the largest is $60$ \\%, which is comparison between $N=128$ clusters in parent galaxy with the steepest and the shallowest mass density profiles.  The important result is that orbital family is attributed to the difference of the mass-loss timescales among star clusters in parent galaxy with different mass density profiles. Even if potential escapers occupy the same phase space among different parent galaxy models, the orbits and the escape time delays may be different by an order of magnitudes. Thus, the difference of mass-loss timescale among different parent galaxies is not represented by orbit-averaged method, such as Fokker-Planck simulations. Orbit-averaged methods use escape criterion based on phase space of cluster stars, for example energy and angular momentum.  We derive formula for the mass-loss timescale of clusters, as seen in equation (\\ref{eq:apply1}). Our formula can explain a property of the population of the galactic globular cluster that their half-mass radii become smaller as their distances from the galactic center become smaller.  \\bigskip"
    },
    "1005/1005.4688_arXiv.txt": {
        "abstract": "We investigate the role of stellar mass in shaping the intrinsic thickness of galaxy discs by determining the probability distribution of apparent axis ratios ($b/a$) for two different samples that probe the faint end of the galaxy luminosity function. We find that the $b/a$ distribution has a characteristic 'U-shape' and identify a limiting mass M$_{*} \\approx 2\\times10^{9}$ M$_{\\odot}$ below which low-mass galaxies start to be systematically thicker. This tendency holds for very faint  ($M_{B} \\sim -8$) dwarfs in the Local Volume, which are essentially spheroidal systems. We argue that galaxy shape is the result of the complex interplay between mass, specific angular momentum and stellar feedback effects. Thus, the increasing importance of turbulent motions in lower mass galaxies leads to the formation of thicker systems, a result supported by the latest hydrodynamical simulations of dwarf galaxy formation and other theoretical expectations. We discuss several implications of this finding, including the formation of bars in faint galaxies, the deprojection of H{\\sc i} line profiles and simulations of environmental effects on the dwarf galaxy population. ",
        "introduction": "Galaxy discs are key structural components in the understanding of galaxy formation. They contain approximately 60\\% of the stellar mass in the Universe \\citep{Driver2007} and are one of the main sites of current star formation activity \\citep{Kennicutt1998}. Moreover, their prominence and appearance form the basis of the \\citet{Hubble1926} morphological sequence.  Despite this enormous importance, understanding the details of disc formation remains extremely challenging. The basic scenario considers that baryons cool and collapse within hierarchically assembled dark matter haloes. A certain amount of their torque-acquired angular momentum is transferred to the baryonic component leading to the formation of a rotationally supported thin disc \\citep{Fall1980}. This relatively simple picture has been able to reproduce many observational properties of disc galaxies --flat rotation curves, Tully-Fisher relation, gas content (e.g., \\citealt{Dalcanton1997,Mo1998,vandenBosch1998,vandenBosch2000})-- but suffers from several shortcomings. On the one hand, detailed N-body simulations showed that when the dissipative effects of gas are not considered, the hierarchical nature of structure formation unavoidably results in disc destruction \\citep{Toth1992}. On the other hand, hydrodynamical simulations of this process systematically produced discs that were too small and too centrally concentrated due to excessive angular momentum exchange between the gas and the dark matter haloes \\citep{Navarro1997}.  A physically-motivated solution for these problems is the inclusion of strong feedback effects from different sources --but specially from star formation and supernovae explosions-- which, when coupled with a cosmic UV field, are able not only to produce realistic discs (e.g., \\citealt{Governato2010}), but also to provide one possible explanation for the missing satellite problem \\citep{Klypin1999}.  This approach, in turn, raises the question of which role does mass play in shaping the properties of discs, as these strong heating mechanisms are expected to produce a greater influence in lower-mass galaxies (\\citealt{Kaufmann2007}, KWB07 hereafter). Indeed, it is well known that galaxies have higher gas mass fractions (e.g., \\citealt{Schombert2001}) and more extended star formation time scales \\citep{Hunter1985} as they are less massive, and it has been proposed that dwarf and disc galaxies are probably two different structural entities \\citep{Schombert2006}.  From an observational point of view, galaxy discs are best described as flattened triaxial ellipsoids with exponential surface brightness profiles \\citep{Freeman1970}. Though they are traditionally considered to be perfectly circular,  it is well known that discs are indeed slightly elliptical, with $b/a \\gtrsim 0.9$ \\citep{Lambas1992,Ryden2004}. They are said to be \\emph{thin} because their vertical to radial axis ratios have long been known to lie in a narrow range $0.15 \\lesssim q_{0} \\lesssim 0.25$ \\citep{Holmberg1950,Sandage1970,Heidmann1972}. However, this thickening is not uniform, as late-type spirals have thinner discs than early-types \\citep{Bottinelli1983,Guthrie1992}. If we neglect these small ellipticity deviations, discs can then be thought of as axisymmetric oblate spheroids with intrinsic thickness $q_{0}$. In that case, the apparent axis ratio of a disc is determined by $b/a = \\sqrt{q_{0}^{2}+(1-q_{0})^{2}\\cos^{2}i}$, where $i$ is the corresponding inclination angle. Interestingly, the apparent axis ratio distribution of a randomly oriented population of such discs peaks at a minimum value of $q_{0}$ --i.e., edge-on galaxies are much more common than any other projection. This property implies that the intrinsic thickness of a disc population can be identified as the minimum value of the distribution of apparent axis ratios.  In this paper we take advantage of this prediction to investigate the range of masses where thin discs exist. For this purpose we study the probability distribution of $b/a$ as a function of galaxy mass (luminosity) for two different samples that probe the faint end of the galaxy luminosity function. In Section\\,2 we develop further on the samples characteristics and the methodology. Section\\,3 presents the main results, which are discussed in Section\\,4 together with additional implications of this investigation. Throughout this Letter we adopt a $\\Lambda$CDM cosmology with $\\Omega_{0} = 0.3$, $\\Omega_{\\Lambda} = 0.7$ and $h_{75} = H_{0}/(75$ km\\,s$^{-1}$\\,Mpc$^{-1})$.    \\section[]{Sample and methodology} Two different galaxy samples have been used throughout this work. The first sample consists of \\emph{all} 9245 galaxies in the SDSS-DR7 \\citep{Abazajian2009} with recession velocities in the $2000 < cz < 6000$ km\\,s$^{-1}$ range. The lower limit was set to avoid strong corrections in galaxy distances due to the Virgo infall velocity field \\citep{Tammann1985}. The upper one ensures that the sample is volume-limited for galaxies brighter than $M_{i} \\approx -16.5$, additionally including 2063 fainter objects down to a limiting magnitude of $M_{i} \\approx -14.5$. We however note that this incompleteness only affects the statistical significance, but not the shape, of the axis ratio probability distribution at the lowest luminosities, as the volume we probe is too narrow ($25 \\lesssim D \\lesssim 80$ Mpc) for any shape evolution to take place. A further concern might arise given the well-known incompleteness of the SDSS imaging for low surface brightness galaxies (LSB; \\citealt{Blanton2005}). However, if a population of LSB discs exists, we will preferentially miss the face-on (high $b/a$) systems instead of the edge-on, higher surface brightness discs. Therefore, the fraction of highly flattened galaxies is in any case a lower limit.  The SDSS imaging provides $ugriz$ magnitudes and several estimates of axis ratios in each band. For this study we have used the 25 mag\\,arcsec$^{-2}$ isophotal axis ratios, which provide a robust measurement of galaxy shapes at their outer regions --a few times the galaxy's effective radius \\citep{Vincent2005}--, where the effects of dust and the presence of bulges are less important. Seeing effects do not play a role either, as the proximity of the sample ensures that essentially all galaxies have minor axis sizes larger than the typical seeing FWHM, with a median of 20 arcsec in the $r$-band. SDSS isophotal $b/a$ can sometimes take unrealistic values in a given photometric band. To avoid these, we have used the median of the three high signal-to-noise $gri$ axis ratio estimates as our reference $b/a$ value, therefore producing a less noisy ($\\sigma_{b/a} \\approx 0.05$) shape estimate for each galaxy. We however verified that our results are independent of the selected filter.   Given that the main purpose of this work is to study the influence of galaxy mass in the formation of discs, we have used the \\citet{Bell2003} M/L ratios with a \\citet{Kroupa2001} initial mass function in the transformation from luminosities to stellar masses.   In order to extend our analysis to even fainter luminosities, we also analyse the catalogue of neighbouring galaxies of \\citet{Karachentsev2004}. This is an all-sky catalogue including 445 galaxies brighter than $M_{B} \\lesssim -8$ with individual distance estimates $D \\leq 10$ Mpc and is expected to be $70\\%$-$80\\%$ complete within 8 Mpc. From this catalogue we used the $B$-band absolute magnitudes and the 25 mag\\,arcsec$^{-2}$ isophotal apparent axis ratios $b/a$ of each individual galaxy. ",
        "conclusions": "Hints of increasing thickness for less massive galaxies are not new (e.g., \\citealt{Heidmann1972,Yoachim2006}), but to our best knowledge we are the first to identify a characteristic mass below which low-mass galaxies start to be systematically thicker. The question then arises naturally. What causes this effect?  One possibility is that environmental effects could influence galaxy shape. While the narrow volume we probe does not include any significant massive cluster --it is just further than Virgo and closer than Coma--, there remains the possibility that low-mass galaxies are tidally affected by more massive companions in group environments, as has been previously observed in the Local Group (\\citealt{Choi2002}; \\citealt{Geha2006a}). To test for this scenario, for each SDSS galaxy we have computed a \\emph{tidal parameter} (e.g., \\citealt{Varela2004}) $\\mbox{T$_{P}$} = \\mbox{max[(M$_{*,P}$/M$_{*,i}$)\\,(A$_{25}$/D$_{P}$)$^{3}$}]$, where M$_{*,i}$ and  A$_{25}$ are its stellar mass and isophotal major axis, and  M$_{*,P}$ and D$_{P}$ are the corresponding stellar mass and projected distance to any potential perturber having a relative velocity $|\\Delta\\mbox{v}| < 1000$ km\\,s$^{-1}$ and a luminosity $M_{i} < -16.5$. This parameter is proportional to the maximum ratio of external to internal forces that act on a galaxy, and takes into account that the effects of a closer, less massive companion can be even more important than those originated by a massive, more distant galaxy. Fig.\\,\\ref{fig:2} shows in grey scale the probability distribution of T$_{P}$ in each mass interval for the SDSS sample, while the different lines indicate the corresponding 1\\%, 10\\%, 50\\%, 90\\% and 99\\% percentiles. There is a clear trend of decreasing T$_{P}$ with decreasing mass, i.e., low-mass galaxies are in general \\emph{less} tidally affected than their more massive counterparts. This result is consistent with the fact that the galaxy luminosity function is dominated \\emph{at all luminosities} by central galaxies \\citep{Cooray2005}. Given that the normalisation of the tidal parameter is somewhat arbitrary, we have used the Hyperleda database \\citep{Paturel2003} to compute the  T$_{P}$ of three dwarf galaxies experiencing different levels of interaction with M31 (see \\citealt{Karachentsev2004}). We find that the fraction of strongly tidally perturbed galaxies like M32 or N205 is only $\\sim$10\\% , while the vast majority of low-mass galaxies are rather isolated systems as WLM.  We therefore conclude that tidal effects are not responsible of the increasing thickness found in fainter galaxies, and suggest that the effect is probably related to the increasing importance of feedback mechanisms in low-mass haloes.  \\begin{figure} \\centering \\includegraphics[width=0.4\\textwidth,clip=true]{sdss_LV_tidalp.eps} \\caption{ Probability distribution of the tidal parameter in bins of stellar mass (grey scale) for the SDSS sample. Note that low-mass galaxies are in general less tidally affected than their more massive counterparts. For reference we show the corresponding values for three local-volume dwarf galaxies experiencing different levels of interaction with M31. } \\label{fig:2} \\end{figure}  KWB07 investigated the effect that a gas temperature floor ($T_{F}$) --as might arise in the presence of a cosmic UV field and/or due to stellar feedback-- has on the properties of low-mass galaxies. Their instructive fig.\\,1 shows that the final morphology of a galaxy is the result of a complex interplay between the turbulent support provided by this $T_{F}$ and the angular momentum support, characterised by the dimensionless spin parameter $\\lambda$. For low-mass haloes, the pressure support radius becomes comparable to or larger than the rotational support radius, and thus galaxies are naturally formed thicker.  \\citet{Schombert2006} was the first to note a separation between dwarf and disc galaxies in the scale-length versus mass plane, in the sense that the former are more diffuse (extended) than the latter. He suggests that the increasing importance of turbulent motions in dwarf galaxies  results in thicker systems, potentially giving rise to the two observed sequences. However, he also points out that 'this scenario provides no explanation for why there are two distinct sequences rather than a smooth transition in structure from an oblate shape to a triaxial one as one progresses to lower galaxy mass'. We believe that the key to this separation resides again in the complex interplay between galaxy mass, heating mechanisms and specific angular momentum.  Going back to KWB07's fig.\\,1, we can identify three different regions. For the most massive haloes, angular momentum support prevails regardless of spin parameter or temperature floor value. In this region, thin discs will form, lying in Schombert's sequence defined by Sc-Sd galaxies. In the case of very low-mass haloes, the opposite is true: pressure support is always dominant even for haloes with the highest spins $\\lambda \\sim 0.1$. Galaxies in this region will be puffier, creating the Im and dIrr sequence in the scale-length vs. mass relation. Finally, for the intermediate mass region KWB07 point out that a range of morphologies is predicted if galaxies have spins within the expected range ($0.01 \\lesssim \\lambda \\lesssim 0.1$; \\citealt{Bullock2001}) and the temperature floor does not vary too much among galaxies. This would naturally explain the co-existence of the two sequences in the narrow $10^{9}$-$10^{10}$ M$_{\\odot}$ mass range: galaxies in the high-spin tail of the distribution would end up in the disc sequence, while the dwarf sequence would be populated by galaxies with low spin parameter.  Evidences for the decreasing importance of ordered motions in favour of turbulent ones come from the high thickness of the neutral gas component in very faint galaxies from the FIGGS sample \\citep{Roychowdhury2010}, or from the shape of H{\\sc i} line profiles from the dwarf sample of \\citet{Geha2006}, where only $18\\%$ of the objects exhibit double-horned profiles (but see \\citealt{Schombert2006}). Furthermore, this fraction increases to $30\\%$ if only edge-on galaxies are considered, suggesting that the shape of these flat objects can be driven by rotation. Indeed, the projected ellipticity of a rotationally supported oblate spheroid with a constant anisotropy parameter is an increasing function of $v/\\sigma$ \\citep{Binney1978}, so that objects with prevailing ordered motions are expected to be more flattened. This picture is supported by fig.\\,8 of \\citet{Geha2006}, where it is shown that the small number of dwarf galaxies with low apparent axis ratios always lie above the best-fit baryonic Tully-Fisher relation \\citep{McGaugh2000}, i.e., their rotation is higher than the mean for their baryonic masses.  How do all these results compare with ours? The minimum of our $b/a$ distribution occurs at a characteristic mass (luminosity) M$_{*} \\approx 2\\times10^{9}$ M$_{\\odot}$ ($M_{i}\\sim-18$). This is roughly the same value as the lower boundary of Schombert's disc sequence, and would thus indicate a limiting mass for thin discs. The baryonic Tully-Fisher relation derived by Geha et al. indicates that galaxies below these luminosities reside in haloes with $V_{c} \\lesssim 80$ km\\,s$^{-1}$ and this suggests that if their thickness is set by turbulent motions, the associated temperature floor must be high (see fig.\\,1 of KWB07).  In order to check if stellar feedback effects can be effectively responsible of the previous trend, we have derived the intrinsic thickness of the dwarf galaxies recently modelled by \\citet{Governato2010}. These hydrodynamical simulations have first produced realistic dwarf galaxies thanks to their high resolution and a detailed treatment of baryonic processes.  Their simulated DG1 has a stellar mass M$_{*} = 4.8\\times10^{8}$ M$_{\\odot}$ and a dimensionless spin parameter $\\lambda = 0.05$. We have measured a 25 mag\\,arcsec$^{-2}$ isophotal axis ratio of $q_{0} = 0.3$ using an $r$-band edge-on image of the dwarf (five-pointed star in Fig.\\,1, left). This value is higher than the minimum thickness we obtain for similar mass galaxies, and we speculate as to whether these flatter objects could reside in haloes with higher spin parameters. Interestingly, the second galaxy simulated by Governato et al. (DG2) --which is slightly less massive, has the same $\\lambda$ and had a quieter merging history than DG1--, is even thicker ($q_{0} = 0.35$, see Fig.\\,1), nicely following the trend with mass we find. We therefore suggest that star formation feedback effects --which can remove low angular momentum material and produce bulgeless dwarfs with shallow central dark matter profiles-- are also responsible of the increasing thickness of low mass galaxies. The shallower potential well certainly allows more turbulent (and therefore also vertical) motions, so stars naturally settle in a thicker disc. This scenario is additionally appealing as it can provide an explanation (cf. KWB07) for the observed high gas mass fractions and low star formation efficiencies of faint systems \\citep{Warren2007,Amorin2009}.  The results presented in this study have several further implications: \\noindent \\begin{itemize} \\item[1)] The existence of a minimum stellar mass for thin discs has a pronounced impact on the formation of bars, as these structures are thought to be formed through instabilities in cold (thin), rotationally supported discs. Therefore, bar formation could be inhibited in galaxies less massive than M$_{*} \\approx 2\\times10^{9}$ M$_{\\odot}$ simply because they are not cold enough. This is consistent with the result recently obtained by \\citet{MendezAbreu2010}, which show that bars are only found in the narrow $10^{9}$-$10^{11}$ M$_{\\odot}$ mass range.%  \\item[2)] It has long been noticed (e.g., \\citealt{Bottinelli1983}) that a varying intrinsic thickness can introduce serious systematic errors in calculated inclinations if they are computed assuming discs have a universal, fixed $q_{0}$ --as it is actually usually done when deprojecting H{\\sc i} line profiles. While errors are only of a few per cent if the galaxy is still an oblate spheroid, the situation becomes more critical if dwarfs are actually better represented as triaxial ellipsoids (e.g., \\citealt{Sung1998}). Moreover, the recent results of \\citet{Roychowdhury2010} indicate that gas discs are \\emph{also} thicker in faint galaxies, further increasing the previous concerns.%  \\item[3)] As already pointed out by KWB07, simulations that investigate the role of environmental effects on the population of dwarf galaxies usually consider a low-mass \\emph{thin} disc as their fiducial model. Given that real dwarf galaxies are intrinsically thicker --and have shallow central dark matter profiles--, they are likely more susceptible to dynamical interactions than most simulations predict.% \\end{itemize}  Interestingly, the evidence of increasing thickness with decreasing mass can be found in several other samples of field late-type dwarfs in the literature (e.g., \\citealt{Eder2000,Geha2006})  as well as for non-nucleated dEs in the Virgo cluster \\citep{Lisker2007}, further suggesting that it might be related to a formation process independent of environment. We believe that future and detailed hydrodynamical simulations of dwarf galaxy formation --probing a variety of masses, spin parameters and merging histories-- should be able to reproduce the probability distribution of apparent axis ratios and provide more quantitative explanations of the observed trends."
    },
    "1005/1005.3899_arXiv.txt": {
        "abstract": "{The equilibrium structure of sunspots depends critically on its magnetic topology and is dominated by magnetic forces. Tension force is one component of the Lorentz force which balances the gradient of magnetic pressure in force-free configurations.} {We employ the tension term of the Lorentz force to clarify the structure of sunspot features like penumbral filaments, umbral light bridges and outer penumbral fine structures.} {We compute vertical component of tension term of Lorentz force over two active regions namely NOAA AR 10933 and NOAA AR 10930 observed on 05 January 2007 and 12 December 2006 respectively. The former is a simple while latter is a complex active region with highly sheared polarity inversion line (PIL). The vector magnetograms used are obtained from Hinode(SOT/SP).} {We find an inhomogeneous distribution of tension with both positive and negative signs in various features of the sunspots. The existence of positive tension at locations of lower field strength and higher inclination is compatible with the uncombed model of the penumbral structure. Positive tension is also seen in umbral light bridges which could be indication of uncombed structure of the light bridge. Likewise, the upward directed tension associated with bipolar regions in the penumbra could be a direct confirmation of the sea serpent model of penumbral structures. Upward directed tension at the PIL of AR 10930 seems to be related to flux emergence. The magnitude of the tension force is greater than the force of gravity in some places, implying a nearly force-free configuration for these sunspot features.} {From our study, magnetic tension emerges as a useful diagnostic of the local equilibrium of the sunspot fine structures.} ",
        "introduction": "The equilibrium structure of sunspots is obviously dominated by magnetic forces. Early researches on this problem dealt mainly with the problem of the global equilibrium of sunspots \\citep{meye77}. The sunspot was modelled as a magnetic flux rope where the lateral force balance was envisaged as a pressure balance between the photospheric plasma pressure and the magnetic + plasma pressure inside the flux rope \\cite[e.g.,][]{chit63}. Since the magnetic field is divergence free, the field lines of the flux rope must bend back into the photosphere resulting in a ``closed'' field topology. The resulting curvature of the field lines produces a magnetic tension that should basically have a downward vertical component. The model of \\cite{meye77} also put constraints on the field line curvature, since the equilibrium becomes unstable when the radius of curvature is shorter than a certain value. With the availability of high resolution magnetograms of sunspots, it has become clear that the earlier models of sunspots might no longer be adequate to explain the dynamical equilibrium of various fine structures seen in the umbra as well as in the penumbra. One simple diagnostic of the vertical equilibrium is the vertical component of the magnetic tension, which can be determined from the lateral gradients of the vector magnetic field. Information about the vertical component of the magnetic tension was already found to be very useful, as e.g. seen in the correlation of low tension force with large magnetic shear in early vector magnetograms \\citep{venk93} measured by the MSFC vector magnetograph \\citep{hagy82}. The magnetic tension measurements at the polarity inversion lines underlying filaments/prominences are like-wise very important since the vanishing of magnetic tension at these highly sheared locations makes the prominence structure extremely vulnerable to dynamical instabilities, via thermal instabilities \\citep{venk90}.  Modern observations of sunspots have revealed the existence of different fine structures. In the umbra we have the umbral dots and light bridges \\cite{sobo89,sobo97,sobo97a}. In the penumbra we have spines (stronger, more vertical field) wrapping around the intraspines (weaker, more horizontal field) \\citep{lites93,borr08}. A few models have also been proposed to explain these structures (e.g., uncombed model of \\cite{sola93}, the gappy model of \\cite{spru06}).  In this paper, we investigate the height variation of the sunspot fine structure using the calculations of the vertical component of the magnetic tension force. The expression for computing the vertical component of tension force is given in Section 2. In Section 3, we describe the data sets used. Section 4 describes the analysis and results. Finally in Section 5 we present our conclusions. ",
        "conclusions": "}  The variation of magnetic field parameters with optical depth can be obtained \\citep{coll94,west01,west01a,beck08} using inversion schemes like SIR \\citep{ruiz92,ruiz94,del96,del03}. With the availability of high quality stokes profiles from the Solar Optical Telescope \\citep{tsun08,suem08,ichi08,shim08} aboard Hinode \\citep{kosu07}, these inversions are now possible with unprecedented quality. In particular, this was used to demonstrate the variation of magnetic field of penumbral filaments as a function of the optical depth \\citep{borr08}, which showed the wrapping of the spine around the intraspine. An inspection of Figure 2 of \\cite{borr08} shows that the magnetic field lines can be inferred to have an upward directed curvature above the intraspine. In the case of our vector magnetograms, the inversions provide the magnetic parameters at a single optical depth in the atmosphere. It has been recognized that the simple inversions based on optical depth independent magnetic parameters are consistent with a weighted mean of the optical depth dependent inversions like SIR \\citep{west98}. We can thus conjecture that our locations of positive tension are manifestations of the upward directed field curvatures that would be seen in SIR inversions. This conjecture must be verified with detailed SIR inversions over the entire sunspot, which will be presented in a future detailed paper. For the present, we can only note that the locations of positive tension in the penumbra appear to be in the regions of lower field strength and larger inclination, which are the properties of the intraspine. We can thus conclude that the lateral variation of the magnetic field, as inferred from simplistic inversions of the stokes profiles, appear to show an upward directed tension, which may be in fact located above the intraspine.  Likewise, the upward directed tension seen in light bridges could well be due to a similar bending of field lines around a more horizontal field \\citep{sankar07}. Finally, the positive tension seen at the edges in bipolar sections of the filamentary penumbral structures could well be a signature of the ``sea-serpent\" \\citep{sain08} where the penumbral fields dip into the photosphere due to hydrodynamic forces caused by convection.  Apart from the sign, the magnitude of the vertical tension also has some very interesting implications. As mentioned earlier, the magnitude of tension attains values upto 10$^{-2}$ dynes/cm$^3$ in some places. This value is comparable to the force of gravity which ranges from 5$\\times10^{-3}$ dynes/cm$^3$ for quiet Sun density to 10$^{-2}$ dynes/cm$^3$ for umbral densities. This must mean that non-magnetic forces alone will not be able to balance this tension force. The only other force that can match this force is the gradient of magnetic pressure. As an example, we can use the values of spine ($\\sim 2000G$) and intraspine ($\\sim 1000G$) of \\cite{borr08} and a scale height $\\Lambda$ of $\\sim$100 km for the vertical variation of field strength, to estimate a downward magnetic pressure gradient ($\\frac{\\Delta B^2}{8\\pi\\Lambda}$) of $\\sim 1.2\\times10^{-2}$ dynes/cm$^3$. This value is of the same order as that of magnetic tension obtained in our analysis. It is well known that configurations where the magnetic tension balances the magnetic pressure gradient are the so-called force-free configurations. Thus, the use of high resolution magnetograms leads to the important result that the existence of upward magnetic tension in sunspot penumbral fine structure could be a manifestation of uncombed magnetic field in {\\it force-free equilibrium}.  Since the observations refer to values at a given optical depth, the resultant magnetogram is not at a single geometrical height. What could be the effect of this corrugation in the height levels? One can estimate that the relative amplitude of tension fluctuations ($\\Delta T/ T$) will be proportional to the relative amplitude in the field variations ($\\Delta B/ B$). Assuming a vertical gradient of magnetic field of 1G/km in the penumbra and a corrugation amplitude of 100 km for the optical depth, we obtain $\\Delta B$ of $\\sim$ 100 G. Knowing that the penumbral fields are $\\sim$ 1000 G, we can estimate the relative field variation as 10\\%, which then is also the estimate for the relative variation in the tension force. The observed spatial variation in tension is much greater than this estimated artefact in tension fluctuations.  \\begin{figure}[h] \\centering \\includegraphics[width=0.48\\textwidth]{f5.eps} \\caption{Contours of tension forces overlaid on the continuum map of NOAA AR 10930 observed on 12 December 2006 at an heliocentric angle of $\\theta$ = 31$^\\circ$. Blue (red) colors show negative (positive) contours of $\\pm1.2, \\pm4, \\pm12$ millidynes/cm$^3$. Polarity inversion line is shown by dashed yellow line. An arrow points towards the center of the solar disk (CSD).} \\label{F4} \\end{figure}   We must remember that the 180 degree ambiguity is not an issue whenever the transverse magnetic vector makes an angle less than 90 degrees with the potential field. In all such cases, the ambiguity is properly resolved by the acute angle method. The NOAA AR 10933 is away from the polarity inversion line and is not highly sheared \\citep{tiw09b,venk09}. We note that the locations of positive tension discussed thus far all have transverse vectors which are not perpendicular to the potential fields. Hence they have no problem as far as the acute angle method is concerned. We do see however that this will be an issue for highly sheared portions of the magnetic field. We normally expect low values of magnetic tension at such highly sheared regions. But in the case of AR 10930, there is strong magnetic tension seen in the highly sheared regions at the polarity inversion line (Figure 3). In such cases, we need to have a careful look at the 180 degree ambiguity resolution before we can decide on the sign of magnetic tension. However, the upward tension could indeed be a signature of the forces that drive flux emergence \\citep{lites97}. This conjecture is consistent with the copious emergence of flux observed on 12 Dec 2006 at the polarity inversion line (PIL) of the AR 10930 \\citep{schr08}. In contrast, the PIL of AR 10930 on 14 December 2006 in Figure A3 (online) shows a reduced value of the tension corresponding to a time of negligible flux emergence. Thus, the study of magnetic tension in various types of sunspot fine structures promises to yield new and exciting information on the equilibrium and dynamics of these structures."
    },
    "1005/1005.1841_arXiv.txt": {
        "abstract": "We  have  investigated the  young  stellar  population  in and  around SFO~38, one  of the massive globules  located in the  northern part of the galactic  \\HII\\ region IC~1396,  using the {\\em Spitzer}  IRAC and MIPS observations  (3.6 to  24 ~\\micron) and  followed up  with ground based  optical photometric and  spectroscopic observations.   Based on the IRAC and MIPS colors and \\ha\\ emission we identify $\\sim$~45 Young Stellar  Objects (Classes 0/I/II)  and 13  probable Pre  Main Sequence candidates.   We  derive the  spectral  types  (mostly  K- and  M-type stars),  effective  temperatures  and  individual  extinction  of  the relatively bright and optically visible Class II objects.  Most of the Class II  objects show variable \\ha\\  emission as well  as optical and near-infrared  photometric  variability  which  confirm  their  ``{\\em youth}\".  Based on optical photometry and theoretical isochrones, we estimate  the spread in  stellar ages  to be  between 1--8~Myr  with a median age of 3 Myr and a mass distribution of 0.3--2.2~$\\msun$ with a median value around 0.5~$\\msun$.  Using the width of the \\ha\\ emission line measured  at 10\\%  peak intensity, we  derive the  mass accretion rates   of   individual   objects   to  be   between   $10^{-10}$   to $10^{-8}$~$\\msun$/yr.  From the  continuum-subtracted \\ha\\ line image, we  find  that the  \\ha\\  emission of  the  globule  is not  spatially symmetric with respect  to the O type ionizing  star HD~206267 and the interstellar  extinction towards  the globule  is also  anomalous.  We clearly  detect  an  enhanced  concentration  of YSOs  closer  to  the southern rim of  SFO~38 and identify an evolutionary  sequence of YSOs from the rim to the dense core of the cloud, with most of the Class II objects located at the bright rim. The YSOs appear to be aligned along two different directions towards the O6.5V type star HD~206267 and the B0V type star HD~206773.  This is consistent with the Radiation Driven Implosion  (RDI)  model for  triggered  star  formation.  Further  the apparent  speed of sequential  star formation  is consistent  with the speed  of propagation  of shocks  in  dense globules  as derived  from numerical simulations of RDI. ",
        "introduction": "Observational and theoretical studies  of star formation over the last decade  have increasingly  strengthened  the idea  that massive  young stars play an important role in triggering the formation of subsequent generations  of  stars.  The  triggers  of  star  formation which  are typically in the form of wind  and radiation from the massive stars or expansion  of  \\HII\\  regions  essentially involve  the  induction  of compression of  a molecular  cloud externally. Among  several proposed theories, two models have gathered sufficient observational support in order to be  regarded as the most plausible  models for triggered star formation. In the first model  known as the {\\em Collect and Collapse} \\citep{elmegreen77,hosokawa06} model, an expanding \\HII\\ region sweeps up material into a dense  bordering layer between the \\HII\\ region and the molecular  cloud. The compressed  shell of dust and  gas undergoes fragmentation   and  gravitational   collapse  to   form   new  stars. Observational  evidence supporting the  ``collect and  collapse\" model includes  detection of  a dense,  fragmented shell  of gas  with newly formed       stars       surrounding       the      \\HII\\       region \\citep{deharveng03,deharveng09,zavagno06,zavagno07}.  The {\\em collect and collapse}  model is however  primarily to explain  the triggered formation  of massive  stars. An  equally likely  model based  on {\\em Radiation  Driven   Implosion  (RDI)},  involves   the  creation  of shock front at the surface  of molecular clouds due to photoevaporation from the surfaces  of molecular clouds exposed to  the UV radiation of nearby massive  young stellar clusters.  The  enhanced inward pressure triggers  the  formation  of  new  protostellar  cores  or  compresses pre-existing ones  to form a  new generation of  stars.  Observational evidence for  the RDI model includes typical  spatial distribution and gradient   in   evolutionary   phases   of   young   stellar   objects \\citep{ikeda08}.   There is  however considerable  ambiguity regarding whether  RDI actually  induces  the gravitational  collapse or  merely exposes the young stars by photoevaporation, since these two scenarios are  observationally  indistinguishable.   Thus  additional  tests  to ascertain  the temporal  feasibility  of shock  waves  from nearby  OB associations  actually   triggering  sequential  star   formation  are required to strengthen the arguments in favor of RDI.  Bright-rimmed clouds  (BRCs) are isolated molecular  clouds located on the  edges of  evolved \\HII\\  regions, which  are found  to  have many signposts of star formation, i.e., {\\em IRAS} point sources, molecular outflows, HH objects  etc. \\citet{sugitani99} and \\citet{sugitani2000} observed a sample  of 89 BRCs and found an  excess in the luminosities and  luminosity-to-cloud ratios  of embedded  {\\em IRAS}  sources when compared with sources in isolated  globules, which is an indication of enhanced star  formation in the BRCs. Young  stellar objects seemingly aligned along the axis towards  the ionizing cluster were detected via near-infrared  (NIR)  imaging of  44  of  these  BRCs. There  is  also evidence for an age gradient, with older stellar objects closer to the OB cluster and  the younger objects well inside  the globules, aligned with the  IRAS sources.  These results are  consistent with sequential formation of stars while the  shock front advances further and further into the  molecular cloud.  Thus  BRCs are ideal objects  for studying and verifying the different models of triggered star formation.   The   \\HII\\    region   IC~1396,    powered   by   the    O6.5V   type \\citep{stickland95} star HD~206267 in the Trumpler~37 cluster, appears to sweep up a molecular  ring of radius 12~pc \\citep{patel1998} and is surrounded by 11 Bright-Rimmed clouds with embedded IRAS point sources \\citep{sugitani94,sugitani91}.   SFO~38, also  known  as IC~1396N,  is located  at  a projected  distance  of  $\\sim$11~pc  to the  north  of HD~206267 with  the bright rim  corresponding to the  ionization front facing the O6.5V  star.  SFO~38 is located in  the Cep~OB2 association at  a  distance  of  750~pc  \\citep{matthews1979}.  More  recent  {\\em Hipparcos}  parallax  based measurements  estimate  the distance  to Cep~OB2 association  to be 615~pc \\citep{dezeeuw99}.  However in order to  make our  results comparable  to most  published results,  in this paper  we  adopt  a  distance  of  750~pc  for  SFO~38.   The  derived luminosities  may therefore  be overestimated  by more  than  40\\% and physical   sizes   by   20\\%.    The   region   is   associated   with IRAS~21391+5802,  a  very   young  intermediate-mass  object,  with  a luminosity of 235~\\lsun\\  \\citep{saraceno96}, which powers an extended bipolar outflow \\citep{sugitani89}.   Based on millimeter observations \\citet{beltran02}  resolved IRAS~21391+5802 into  an intermediate-mass source  named  BIMA~2  surrounded  by  two less  massive  and  smaller objects,  BIMA~1  and  BIMA~3.   \\cite{valdettaro2005}  detected  {\\em $H_{2}O$}  maser  emission  at   2.2~GHz  towards  SFO~38  which  is consistent  with  an  intermediate-mass object.   \\citet{neri07}  used still    higher   angular   resolution    millimeter   interferometric observations  to reveal  that the  intermediate-mass  protostar BIMA~2 itself  consists  of  multiple  compact  sources.   The  gas  emission surrounding   IRAS~21391+5802  traces  different   molecular  outflows \\citep{codella01,beltran02,beltran04}.  NIR images  of the region have also  revealed  a  number   of  small  scale  molecular  hydrogen  and Herbig-Haro                         (HH)                         flows \\citep{nisini01,sugitani02,reipurth03,caratti06,beltran09}.       These observations confirm the on going  star formation in the dense core of SFO~38.   \\cite{getman07}  identified  the  mid-infrared (MIR)  ({\\em  Spitzer}) counterparts of  X-ray sources detected in  {\\em Chandra} observations and found that  the young stellar objects (YSOs)  are oriented toward the ionizing star  HD~206267 and show an age  gradient consistent with the RDI model for triggered star formation.  Recently \\cite{beltran09} studied  the  YSO  population  of  SFO~38  by  obtaining  deep  J,  H, K$^\\prime$   broadband  images   and   deep  high-angular   resolution observations  in the H$_2$  narrowband filter.   \\citet{beltran09} did not detect any clear NIR excess close to the rim, which is in contrast to   the  age   1~Myr   of  the   stellar   population  concluded   by \\citet{getman07}.  \\citet{beltran09}  suggest that the  YSOs closer to HD~206267 could  actually be  younger than the  Class II  objects they appear  to be  since  their circumstellar  environment were  disrupted completely   by    intense   UV   radiation   field    from   the   OB association. \\citeauthor{beltran09}  thus suggest that  in general the apparent  age  sequence  seen   close  to  OB  associations  need  not necessarily be  the actual  evolutionary sequence and  in the  case of SFO\\,38 there is no concrete  evidence of star formation by the nearby OB association.  In this paper  we address the controversial issue  of the distribution of YSOs  by making use of  the extremely sensitive  {\\em Spitzer} IRAC (3.6 to 8.0~\\micron) \\& MIPS  (24~\\micron) data in order to probe {\\em all} embedded YSOs and protostars  in SFO~38.  This thus extends the work     by    \\citet{getman07},     who     had    considered     the 3.6--5.8~\\micron\\ characteristics  only of those  MIR sources from  which  X-ray emission  had  been  detected.   Our approach  also complements the  work by \\cite{beltran09} which  makes use of only the  NIR observations to decipher  any age gradient  of the YSOs. In  this  work we  shall  show  that most  of  the  Class 0/I  sources identified  in  the  MIR   are  either  barely  detected  (with  large uncertainties in fluxes) or not detected  at all in the NIR and in the X-rays.  In this paper  we have completed the multiwavelength overview of SFO~38  by obtaining optical  broadband images in {\\em  BVRI} bands and narrow band \\ha\\ image to understand the detailed structure of the bright-rim. Further in order to confirm the ``{\\em youth}\" of the YSOs identified  in  the  MIR  we   have  made  use  of  medium  resolution spectroscopy  of the  optically visible  YSOs in  \\ha\\  line emission. Observation details of  { \\em Spitzer} and optical  {\\em BVRI} imaging and medium resolution  spectroscopy are presented in \\S~\\ref{sec_obs}. We describe the  selection of YSO candidates based  on MIR color-color diagram  (CCD)  and the  detailed  analysis  of  optical--MIR data  in section \\S~\\ref{sec_results}.  In \\S~\\ref{sec_dis} we discuss the probable  star formation scenarios  and in \\S~\\ref{sec_conclusion} we summarize the results. ",
        "conclusions": "}  We  have presented  multiwavelength photometric  study of  SFO~38 from optical {\\em BVRI}  to {\\em Spitzer} IRAC and  MIPS observations along with optical  spectroscopy of the  selected objects. We  summarize our work as follows.  \\begin{itemize} \\item  A  total  of  40  YSOs (Classes  0/I/II)  and  13  YSO candidates are  identified based on  MIR color indices and  we also confirm \\ha\\  emission from  2 YSOs which  are not known  in the literature.  We  further  identify   4  additional  YSOs  based  on \\ha\\ emission and thus we find 44 YSOs in and around SFO~38.  \\item \\ha\\ emission line YSOs are  mostly K--M type stars with an age spread of  1-8 Myr and mass range  of .3-2.2 \\msun\\.  Some of the YSOs show  photometric variation in  optical and NIR  bands and also variable \\ha\\ emission  in the medium resolution spectra. Mass accretion rates estimated form broad  \\ha\\ line profiles are of the order of 10$^{-8}$ to 10 $^{-10}$ \\msun\\ yr$^{-1}$.  \\item Mass, luminosity  and age for the different  components of the protostellar cluster at IRAS 21391+5802 were derived by SED fitting.  \\item   Continuum  subtracted   \\ha\\  line   image   show  asymmetric \\ha\\  emission at  the  bright-rim. Two   OB  type stars e.g.  HD  206267 (O6.5 ) and HD  206773 (B0V) are proposed as the potential ionizing sources.  \\item Class II  to Class 0/I objects are distributed  from the rim to the  core part  of  the  globule respectively.  It  is possible  to identify  at least  two different  axes  of elongation  of the  YSO distribution. The spatial gap  between the two different classes of YSOs is  consistent with the difference  in the arrival  time of a shockwave  propapagating  into  the globule.   The  spatio-temporal gradient in the distribution of  YSOs along two different axes that are parallel to one of  the either ionizing star indicate triggered star formation due to the Radiation Driven Implosion model.  \\end{itemize}  This work  is based  [in part] on  observations made with  the Spitzer Space Telescope,  which is operated by the  Jet Propulsion Laboratory, California Institute  of Technology under a contract  with NASA.  This research  has made  use  of the  SIMBAD  data base,  operated at  CDS, Strasbourg, France.   This publication makes use of  the data products from  the  2MASS,  which is  a  joint  project  of the  University  of Massachusetts    and   the    Infrared    Processing   and    Analysis Center/California  Institute  of Technology,  funded  by the  National Aeronautics  and   Space  Administration  and   the  National  Science Foundation.  RC would  like to thank Dr.  P. S.   Parihar and Ramya S. for useful discussions. \\begin{figure*}[ht] \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0,width=18.0cm,angle=0]{fig/color-com.eps}} \\caption{IRAC-MIPS color-composite image of SFO 38 using 3.6~\\micron\\ (Blue), 8~\\micron\\ (Green) and 24~\\micron\\ (Red). The image is centered at $\\alpha_{2000}$  = 21$^h$40$^m$42$^s$ and $\\delta_{2000}$ = 58\\arcdeg16\\arcmin10\\arcsec and extends over 5\\arcmin$\\times$5\\arcmin\\ ($\\alpha\\times\\delta$).} \\label{fig_iracyso} \\end{center} \\end{figure*}   \\begin{figure*} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics{./fig/ha-mir-color.eps}} \\caption{Continuum subtracted \\ha\\ emission line image of SFO 38. Class 0/I to Class II YSOs are overplotted with different symbols. The directions towards HD 206267 and HD 206773 are marked with two different {\\em arrows}. \\label{fig_ysomap}} \\end{center} \\end{figure*}  \\begin{figure*} \\begin{minipage}{5.5cm} \\includegraphics[width=5.4cm]{fig/irac-cc2.eps} \\end{minipage} \\begin{minipage}{5.5cm} \\includegraphics[width=5.4cm]{fig/irac-cc1.eps} \\end{minipage} \\begin{minipage}{5.5cm} \\includegraphics[width=5.4cm]{fig/near-irac-CCD.eps} \\end{minipage} \\caption{Color-color  diagrams  for   all  MIR   sources. Approximate classification  criteria adopted from \\citet{megeath04}, \\citet{fang09} and \\citet{muzerolle04} are  shown in the two panels. (a) MIR color-color  diagram based on {\\em Spitzer}  3-band IRAC and MIPS  photometry and  (b) 4  band IRAC  color-color  diagram. Dashed lines  in the two panels are  taken from  \\citet{hartman05}.  Reddening vector  corresponding to  the extinction  laws given  by  the fitted function from \\citet{indebetouw05} is shown in {\\em (b)}. The dashed lines [3.6]--[4.5] = 0.7, [4.5]--[5.8]  = 0.7 and [5.8]--[8.0] = 0.7 discriminate  Class   II  sources   from  Class  I/0   sources,  and [3.6]--[4.5]  =  0.2, [4.5]--[5.8]  =  0.2  and  [5.8]--[8.0] =  0.2 discriminate Class III from Class II and Transitional Disk sources. (c) NIR and {\\em Spitzer} MIR color-color diagram. {\\em Solid} line is the reddening vector due to interstellar extinction \\citep{tapia81}. Sources to the right of the reddening vector are YSOs with NIR excess. \\label{fig_iracmipscol}} \\end{figure*}   \\begin{figure*} \\begin{center} \\resizebox{15cm}{!}{\\includegraphics[angle=0,width=15cm,angle=0]{./fig/specplot.eps}} \\caption{Sample spectra of YSOs in SFO 38 obtained with {\\em HFOSC } instrument. The strong features e.g. \\ha\\ emission ($\\lambda 6563$\\AA\\ ), CaII infrared triplet emission , OI line at $\\lambda 7773$ \\AA\\ and [OI] line at $\\lambda 6300$ \\AA\\ are marked in the spectra. Atmospheric features are indicated with $\\earth$. Two spectra of MIR-29 taken at two different epochs show variable \\ha\\ emission (see text for details).} \\label{fig_specplot} \\end{center} \\end{figure*}   \\begin{figure*}[ht] \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0,width=18cm,angle=0]{fig/VI-cm.eps}} \\caption{Extinction corrected V vs. V-I Color-Magnitude diagram of YSOs near SFO 38. \\ha\\ emission YSOs are shown by {\\em filled circles}. The ZAMS and 0.6, 1, 5, 10 Myr isochrones of \\cite{siesspms} are shown with the evolutionary tracks for masses from 0.3 to 2.2 \\msun, at an adopted distance of 750~pc. \\label{fig_vicm}} \\end{center} \\end{figure*}    \\begin{figure*} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics{./fig/sed-mir.eps}} \\caption{SED  fits for  the PMS  objects using  axisymmetric radiation transfer models.   The filled  circles indicate the  measured fluxes and  uncertainties. The  filled  triangles correspond  to the  upper limits  of flux  at these  wavelengths.  The  continuous  black line represents  the  best-fitting SED, and the  dashed line  shows the stellar photosphere corresponding to  the central source of the best fitting model,  in the absence of circumstellar  dust (but including interstellar extinction).} \\label{fig_sedfits} \\end{center} \\end{figure*}  \\newpage \\setcounter{table}{0} \\begin{deluxetable}{lcclllrrrrrrl} \\tabletypesize{\\tiny} \\rotate \\tablecaption{Results of Near- \\& Mid-IR photometry of the Mid-infrared sources in SFO 38 \\label{tab_mirsrc}} \\tablewidth{0pt} \\tablehead{ \\colhead{Source} & \\colhead{$\\alpha_{2000}$}& \\colhead{$\\delta_{2000}$}& \\colhead{$J$}& \\colhead{$H$}& \\colhead{$K_{\\rm s}$}& \\colhead{$F_{3.6}$}& \\colhead{$F_{4.5}$}& \\colhead{$F_{5.8}$}& \\colhead{$F_{8.0}$} & \\colhead{$F_{24}$}& \\colhead{Class} & \\colhead{Comments}\\\\ SFO38 & &&&&&mJy&mJy&mJy&mJy&mJy&\\\\} \\startdata MIR-1 & 21:40:24.97 &  58:16:10.5 &      14.48$\\pm$0.04  &    13.72$\\pm$ 0.05  & 13.42$\\pm$ 0.05 &     1.67$\\pm$0.01&     1.10$\\pm$0.01& 0.96$\\pm$0.02& 1.03$\\pm$0.02&  \\ldots & Trans. & {\\em $JHK_s$}\\tablenotemark{c}\\\\ MIR-2 & 21:40:25.39 &  58:15:47.3 &  15.35$\\pm$0.06  &    14.51$\\pm$ 0.07  & 14.15$\\pm$ 0.07 &     0.94$\\pm$0.01&     0.63$\\pm$0.01& 0.68$\\pm$0.02& 0.98$\\pm$0.02&  \\ldots & Trans. & \\ldots\\\\ MIR-3 & 21:40:25.84 &  58:15:26.1 &   15.62$\\pm$0.07  &    14.80$\\pm$ 0.07  &    14.62$\\pm$ 0.10 &     0.52$\\pm$0.01&     0.31$\\pm$0.00& 0.23$\\pm$0.02&     0.07$\\pm$0.02&  \\ldots &III/S & \\ldots\\\\ MIR-4 & 21:40:26.84 &  58:14:51.8  &   16.26$\\pm$0.10  &    15.67$\\pm$ 0.15  &    15.06$\\pm$ 0.15 &     0.32$\\pm$0.00&     0.20$\\pm$0.00& 0.17$\\pm$0.02&     0.11$\\pm$0.02&  \\ldots &III/S  & \\ldots\\\\ MIR-5 & 21:40:27.29 &  58:14:21.14 & 14.30$\\pm$0.04  &    13.30$\\pm$ 0.04  &    12.88$\\pm$ 0.04 &     6.22$\\pm$0.05& \\ldots& \\ldots& \\ldots& \\ldots & II & \\ha, X-ray\\\\ MIR-6 & 21:40:27.38 &  58:16:21.3 &      \\ldots  &     \\ldots  &     \\ldots &     0.07$\\pm$0.00&     0.03$\\pm$0.00&     0.51$\\pm$0.02& 2.00$\\pm$0.03&  \\ldots &I/II\\\\ MIR-7 & 21:40:27.97 &  58:15:14.1 &    14.51$\\pm$0.04  &    13.41$\\pm$ 0.04  & 12.94$\\pm$ 0.04 &     3.89$\\pm$0.02&     3.52$\\pm$0.02& 3.44$\\pm$0.03& 4.33$\\pm$0.03&  \\ldots & II & \\ha, {\\em J}\\\\ MIR-8 & 21:40:28.35 &  58:16:42.3 &     11.59$\\pm$0.03  &    10.78$\\pm$ 0.03  &    10.43$\\pm$ 0.02 &    23.62$\\pm$0.07&    14.18$\\pm$0.04& 10.94$\\pm$0.06&     6.63$\\pm$0.04&  \\ldots & III/S\\\\ MIR-9 & 21:40:30.63 &  58:15:00.0 &     15.31$\\pm$0.05  &    14.90$\\pm$ 0.07  &    14.70$\\pm$ 0.11 &     0.49$\\pm$0.01&     0.32$\\pm$0.01& 0.28$\\pm$0.01&     0.22$\\pm$0.02&  \\ldots & Trans.\\\\ MIR-10 & 21:40:31.00 &  58:15:09.2 &    13.44$\\pm$0.03  &    12.97$\\pm$ 0.04  &    12.79$\\pm$ 0.04 &     2.61$\\pm$0.02&     1.63$\\pm$0.01& 1.29$\\pm$0.02& \\ldots&  \\ldots & III/S\\\\ MIR-11 & 21:40:31.56 &  58:17:55.1 &     14.03$\\pm$0.03  &    12.89$\\pm$ 0.03  &    12.39$\\pm$ 0.03 &     6.93$\\pm$0.03&     6.93$\\pm$0.03& 9.15$\\pm$0.06&    15.17$\\pm$0.12&     20.46$\\pm$0.08 & II & \\ha, X-ray, {\\em $JHK_s$}\\\\ MIR-12 & 21:40:31.97 &  58:14:56.0 &    16.72$\\pm$0.16 &    15.78$\\pm$ 0.15  &    15.41$\\pm$ 0.21 &     0.41$\\pm$0.00& 0.28$\\pm$0.00&     0.26$\\pm$0.01& \\ldots&  \\ldots &III/S \\\\ MIR-13 & 21:40:32.23 &  58:16:53.6 &    14.13$\\pm$0.03  &    12.44$\\pm$ 0.04  &    11.71$\\pm$ 0.03 &    10.82$\\pm$0.04&     6.39$\\pm$0.02& 6.54$\\pm$0.05&     7.36$\\pm$0.05&  \\ldots & Trans.\\\\ MIR-14 & 21:40:32.46 &  58:13:47.2 &     13.40$\\pm$0.03  &    12.83$\\pm$ 0.03  &    12.63$\\pm$ 0.03 &     2.78$\\pm$0.03& \\ldots& 1.30$\\pm$0.05& \\ldots&  \\ldots& III/S \\\\ MIR-15 & 21:40:32.71 &  58:17:23.5 &     12.19$\\pm$0.03  &    11.55$\\pm$ 0.03  &    11.36$\\pm$ 0.03 &     9.21$\\pm$0.04&     5.50$\\pm$0.02& 4.04$\\pm$0.04&     1.71$\\pm$0.04&  \\ldots& III/S \\\\ MIR-16\\tablenotemark{b} & 21:40:32.84 &  58:15:47.8 & 17.42$\\pm$0.01& 16.61$\\pm$0.01 & 16.09$\\pm$0.01 &     1.42$\\pm$0.01&     0.73$\\pm$0.01&     9.50$\\pm$0.05& 32.10$\\pm$0.11&  \\ldots & I/II\\\\ MIR-17 & 21:40:33.64 &  58:14:44.6 &  16.54$\\pm$0.14  &    15.81$\\pm$ 0.18  &    15.07$\\pm$ 0.00 &     0.21$\\pm$0.00&     0.13$\\pm$0.00& 0.11$\\pm$0.01&     0.16$\\pm$0.02&  \\ldots & Galaxy\\\\ MIR-18 & 21:40:33.69 &  58:14:59.5 &     13.06$\\pm$0.03  &    12.09$\\pm$ 0.03  &    11.75$\\pm$ 0.02 &     8.13$\\pm$0.03&     4.82$\\pm$0.02& 4.72$\\pm$0.04&     5.60$\\pm$0.03&  \\ldots & Trans.\\\\ MIR-19 & 21:40:34.05 &  58:18:08.1 &  15.70$\\pm$0.06  &    13.40$\\pm$ 0.04  &    12.30$\\pm$ 0.02 &     7.41$\\pm$0.03&     5.16$\\pm$0.02& 4.95$\\pm$0.04&     4.03$\\pm$0.04&  \\ldots & Trans.\\\\ MIR-20 & 21:40:34.30 &  58:18:23.6 &  \\ldots  &     \\ldots  & \\ldots &     0.31$\\pm$0.00&     0.84$\\pm$0.01&     1.03$\\pm$0.02& 1.24$\\pm$0.04&  \\ldots& H$_2$ \\\\ MIR-21 & 21:40:34.59 &  58:16:11.4 &  \\ldots  &     \\ldots  &     \\ldots &     0.08$\\pm$0.00&     0.05$\\pm$0.00&     1.33$\\pm$0.03& \\ldots&  \\ldots & III/S\\\\ MIR-22\\tablenotemark{b} & 21:40:34.76 &  58:15:20.9 &   17.70$\\pm$0.01 &  16.68$\\pm$0.01 &  16.04$\\pm$0.01 &     0.25$\\pm$0.01&     0.34$\\pm$0.01& \\ldots& \\ldots&  \\ldots& III/S \\\\ MIR-23 & 21:40:34.79 &  58:16:10.8 &  \\ldots  & \\ldots  &     \\ldots & \\ldots& \\ldots& \\ldots&     4.47$\\pm$0.05& \\ldots& \\ldots \\\\ MIR-24 & 21:40:35.02 &  58:18:22.2 &   16.48$\\pm$\\ldots & 15.72$\\pm$\\ldots  & 14.90$\\pm$ 0.21  &         0.62$\\pm$0.01& 2.68$\\pm$0.01& 1.84$\\pm$0.03&     1.42$\\pm$0.03&  \\ldots & H$_2$ \\\\ MIR-25\\tablenotemark{b} & 21:40:35.11 &  58:15:30.9 &   18.21$\\pm$0.02 &    16.85$\\pm$0.01 & 16.16$\\pm$0.02 &     0.92$\\pm$0.01&     0.48$\\pm$0.01&     5.32$\\pm$0.04& 18.26$\\pm$0.08&  \\ldots &I/II \\\\ MIR-26 & 21:40:35.20 &  58:14:05.0 &     13.11$\\pm$0.03  &    12.61$\\pm$ 0.03  &    12.45$\\pm$ 0.03 &     3.55$\\pm$0.02&     2.40$\\pm$0.01& 1.59$\\pm$0.02&     0.84$\\pm$0.02&  \\ldots & III/S \\\\ MIR-27 & 21:40:35.68 &  58:18:21.3 &    16.55$\\pm$\\ldots & 15.73$\\pm$\\ldots  & 14.79$\\pm$ 0.17  &    0.77$\\pm$0.01& 3.61$\\pm$0.02& 2.20$\\pm$0.03&     1.91$\\pm$0.03&  \\ldots & H$_2$\\\\ MIR-28 & 21:40:36.46 &  58:15:23.6 &   15.97$\\pm$0.10  &    14.73$\\pm$ 0.08  &    14.04$\\pm$ 0.07 &     2.27$\\pm$0.01&     1.43$\\pm$0.01& 6.68$\\pm$0.04&    19.47$\\pm$0.08&  \\ldots & I/II\\\\ MIR-29 & 21:40:36.52 &  58:13:45.8 &  13.51$\\pm$0.02  &    12.58$\\pm$ 0.03  &    12.24$\\pm$ 0.03 &     4.69$\\pm$0.02&     3.14$\\pm$0.01& 2.17$\\pm$0.03&     2.64$\\pm$0.03&     44.29$\\pm$0.09  & Trans. & \\ha, X-ray\\\\ MIR-30 & 21:40:36.65 &  58:18:23.0 &    17.12$\\pm$\\ldots & 15.18$\\pm$0.10 & 13.93$\\pm$0.07 &     1.84$\\pm$0.01& 1.55$\\pm$0.01& 1.42$\\pm$0.03&     1.21$\\pm$0.03&  \\ldots & II \\\\ MIR-31 & 21:40:36.88 &  58:14:37.8 &    11.90$\\pm$0.02  &    10.89$\\pm$ 0.03  &    10.23$\\pm$ 0.02 &    54.30$\\pm$0.14&    48.79$\\pm$0.10& 38.25$\\pm$0.12&    34.45$\\pm$0.10&     39.36$\\pm$0.07& II  & \\ha, X-ray\\\\ MIR-32 & 21:40:37.02 &  58:15:02.9 &     14.27$\\pm$0.03  &    13.28$\\pm$ 0.04  &    12.82$\\pm$ 0.03 &     4.01$\\pm$0.02&     3.07$\\pm$0.01& 2.86$\\pm$0.03&     3.87$\\pm$0.04&  \\ldots& II  & \\ha, {\\em $K_s$}\\\\ MIR-33 & 21:40:37.16 &  58:15:46.6 &  17.01$\\pm$\\ldots & 15.80$\\pm$0.16 & 14.55$\\pm$0.10  &    1.33$\\pm$0.01&     1.28$\\pm$0.01& 0.71$\\pm$0.03& \\ldots&  \\ldots & III/S\\\\ MIR-34\\tablenotemark{b} & 21:40:37.54 &  58:15:39.3 &   18.07$\\pm$0.02 &    16.98$\\pm$0.01 &  16.24$\\pm$0.02 &     0.49$\\pm$0.01&     0.77$\\pm$0.01&     1.05$\\pm$0.03& 1.63$\\pm$0.04&     17.15$\\pm$0.06 & 0/I\\\\ MIR-35 & 21:40:39.05 &  58:16:03.3 &    \\ldots  &     \\ldots  & \\ldots &     0.41$\\pm$0.01&     1.04$\\pm$0.01&     2.19$\\pm$0.03& \\ldots& \\ldots & H$_2$ \\\\ MIR-36\\tablenotemark{a} & 21:40:39.60 &  58:16:09.0    & \\ldots & $>$17.6 & 15.14$\\pm$0.16  &     10.47$\\pm$0.03&    36.11$\\pm$0.08&    59.09$\\pm$0.17& 82.89$\\pm$0.22&    239.10$\\pm$0.10 & 0/I & X-ray\\\\ MIR-37 & 21:40:39.80 &  58:18:24.9 &   \\ldots  &     \\ldots  & \\ldots &     0.53$\\pm$0.01&     0.76$\\pm$0.01&     1.15$\\pm$0.02& 1.39$\\pm$0.03&  \\ldots& 0/I \\\\ MIR-38 & 21:40:39.84 &  58:18:34.8 &   18.29$\\pm$\\ldots &  15.30$\\pm$0.12 & 13.36$\\pm$0.04  &      6.88$\\pm$0.03& 8.19$\\pm$0.03& 9.95$\\pm$0.06&    13.49$\\pm$0.06&     30.69$\\pm$0.07 & II& X-ray \\\\ MIR-39 & 21:40:40.15 &  58:16:18.2 &   \\ldots  &     \\ldots  & \\ldots &     1.36$\\pm$0.01&     2.65$\\pm$0.01&     3.24$\\pm$0.04& 4.25$\\pm$0.05&  \\ldots & 0/I  & \\\\ MIR-40 & 21:40:40.43 &  58:18:26.1 &   18.28$\\pm$\\ldots & 15.64$\\pm$0.15 & 13.85 $\\pm$0.06 &     2.35$\\pm$0.01&     2.05$\\pm$0.01& 1.64$\\pm$0.03&     0.76$\\pm$0.02&  \\ldots & III/S\\\\ MIR-41 & 21:40:40.96 &  58:16:12.0 &   17.55$\\pm$\\ldots & 15.89$\\pm$\\ldots  & 13.59$\\pm$ 0.06  &    13.04$\\pm$0.04& 23.41$\\pm$0.06& 29.43$\\pm$0.10&    42.25$\\pm$0.15&  \\ldots & 0/I \\\\ MIR-42 & 21:40:40.98 &  58:16:52.4 &    17.93$\\pm$\\ldots & 15.34$\\pm$0.14 & 13.93$\\pm$0.06 &     1.41$\\pm$0.01&     1.04$\\pm$0.01& 0.83$\\pm$0.02&     0.37$\\pm$0.03&  \\ldots& III/S \\\\ MIR-43 & 21:40:41.07 &  58:13:58.9 &      12.96$\\pm$0.02  &    12.08$\\pm$ 0.03  &    11.77$\\pm$ 0.02 &     6.80$\\pm$0.02&     4.58$\\pm$0.02& 3.19$\\pm$0.03&     1.84$\\pm$0.03&  \\ldots& III/S  & \\ha, X-ray\\\\ MIR-44 & 21:40:41.10 &  58:17:54.1 &      14.20$\\pm$0.04  &    13.46$\\pm$ 0.04  &    13.21$\\pm$ 0.03 &     1.89$\\pm$0.01&     1.14$\\pm$0.01& 0.73$\\pm$0.02&     0.37$\\pm$0.02&  \\ldots& III/S \\\\ MIR-45 & 21:40:41.14 &  58:15:11.2 &    12.97$\\pm$0.03  &    11.61$\\pm$ 0.04  &    10.68$\\pm$ 0.02 &    66.38$\\pm$0.17&    65.98$\\pm$0.14& 81.84$\\pm$0.21&   124.60$\\pm$0.31&    194.40$\\pm$0.09& II  & \\ha, X-ray\\\\ MIR-46 & 21:40:41.27 &  58:17:14.2 &  \\ldots  &     \\ldots  &     \\ldots &     0.06$\\pm$0.00&     0.03$\\pm$0.00&     0.65$\\pm$0.02& 2.61$\\pm$0.03&  \\ldots & I/II\\\\ MIR-47 & 21:40:41.40 &  58:14:17.5 &   15.48$\\pm$\\ldots & 15.50$\\pm$0.12 & 14.63$\\pm$0.00 &     0.34$\\pm$0.00&     0.21$\\pm$0.00& 0.16$\\pm$0.01&     0.10$\\pm$0.02&  \\ldots& III/S \\\\ MIR-48 & 21:40:41.43 &  58:16:37.8 &    18.30$\\pm$\\ldots & 15.60$\\pm$0.15 &  12.79$\\pm$0.03 &  51.81$\\pm$0.14& 109.40$\\pm$0.20& 158.50$\\pm$0.34&   199.60$\\pm$0.44& 1362.00$\\pm$0.26 & 0/I\\\\ MIR-49 & 21:40:41.54 &  58:14:25.5 &    13.65$\\pm$0.03  &    12.62$\\pm$ 0.03  &    12.17$\\pm$ 0.03 &     8.70$\\pm$0.03&     7.38$\\pm$0.03& 6.65$\\pm$0.04&     6.42$\\pm$0.04&  \\ldots & II & \\ha, X-ray\\\\ MIR-50 & 21:40:41.71 &  58:16:12.6 &     \\ldots  &     \\ldots  & \\ldots &    10.84$\\pm$0.04&    78.01$\\pm$0.16&   128.00$\\pm$0.29& 137.00$\\pm$0.34&   1871.00$\\pm$0.31 & 0/I & X-ray\\\\ MIR-51 & 21:40:41.89 &  58:15:23.0 &  15.68$\\pm$0.09  &    14.30$\\pm$ 0.06  &    13.65$\\pm$ 0.05 &     2.88$\\pm$0.02&     2.30$\\pm$0.01& 3.75$\\pm$0.04&     8.33$\\pm$0.06&  \\ldots & I/II & X-ray, {\\em $JHK_s$}\\\\ MIR-52 & 21:40:41.96 &  58:13:49.1 &     15.50$\\pm$0.06  &    15.00$\\pm$ 0.09  &    14.98$\\pm$ 0.14 &     0.40$\\pm$0.00&     0.27$\\pm$0.00& 0.19$\\pm$0.01&     0.10$\\pm$0.02&  \\ldots & III/S\\\\ MIR-53 & 21:40:42.35 &  58:18:40.1 &   \\ldots  &     \\ldots  &     \\ldots &     0.48$\\pm$0.00&     2.16$\\pm$0.01&     1.99$\\pm$0.03& 1.26$\\pm$0.02&  \\ldots & H$_2$ \\\\ MIR-54 & 21:40:42.77 &  58:16:01.1 &    \\ldots  &     \\ldots  & \\ldots &    20.16$\\pm$0.06&    75.86$\\pm$0.15&   113.50$\\pm$0.27& 117.60$\\pm$0.31&   2484.00$\\pm$0.38 & I  & X-ray\\\\ MIR-55\\tablenotemark{a} & 21:40:43.32 &  58:17:16.5 &    \\ldots &     $>$17.6  & 15.31$\\pm$0.11 &     1.65$\\pm$0.01&     2.23$\\pm$0.01&     2.72$\\pm$0.03& 4.31$\\pm$0.03&    145.40$\\pm$0.08  & 0/I\\\\ MIR-56 & 21:40:43.34 &  58:18:42.6 &  16.32$\\pm$0.13  &    14.42$\\pm$ 0.05  &    13.55$\\pm$ 0.04 &     1.93$\\pm$0.01&     1.39$\\pm$0.01& 0.91$\\pm$0.02&     0.26$\\pm$0.03&  \\ldots & III/S\\\\ MIR-57 & 21:40:43.62 &  58:16:18.8 &   17.89$\\pm$\\ldots & 16.09$\\pm$\\ldots  & 13.51$\\pm$ 0.07  &    33.21$\\pm$0.09& 68.55$\\pm$0.14& 82.72$\\pm$0.21&    91.38$\\pm$0.24&  \\ldots & 0/I & X-ray \\\\ MIR-58 & 21:40:43.65 &  58:14:18.4 &   15.55$\\pm$0.08  &    15.05$\\pm$ 0.10  &    14.89$\\pm$ 0.14 &     0.41$\\pm$0.00&     0.28$\\pm$0.00& 0.19$\\pm$0.01&     0.14$\\pm$0.02&  \\ldots& Trans. \\\\ MIR-59 & 21:40:44.23 &  58:16:52.2 &  \\ldots  &     \\ldots  & \\ldots &     1.36$\\pm$0.01&     4.43$\\pm$0.02&     9.13$\\pm$0.05& 10.52$\\pm$0.06&    100.90$\\pm$0.07 & 0/I \\\\ MIR-60 & 21:40:44.26 &  58:14:12.4 &   15.57$\\pm$0.06  &    15.00$\\pm$ 0.09  &    14.69$\\pm$ 0.12 &     0.45$\\pm$0.00&     0.30$\\pm$0.00& 0.22$\\pm$0.01&     0.09$\\pm$0.02&  \\ldots& III/S \\\\ MIR-61 & 21:40:44.31 &  58:15:13.2 &   16.05$\\pm$0.10  &    14.59$\\pm$ 0.07  &    13.60$\\pm$ 0.05 &     3.78$\\pm$0.02&     4.10$\\pm$0.02& 7.89$\\pm$0.05&    15.74$\\pm$0.08&  \\ldots & I/II & X-ray, {\\em $JHK_s$}\\\\ MIR-62 & 21:40:44.64 &  58:17:04.9 &   17.83$\\pm$\\ldots & 15.75$\\pm$0.15  & 14.62$\\pm$0.10 &     2.16$\\pm$0.01& 2.38$\\pm$0.01& 2.51$\\pm$0.03&     3.06$\\pm$0.03&  \\ldots & II \\\\ MIR-63 & 21:40:44.82 &  58:16:05.0 &   15.86$\\pm$\\ldots & 14.28$\\pm$0.06  & 12.89$\\pm$0.05 &     7.08$\\pm$0.03& 7.17$\\pm$0.03& 6.83$\\pm$0.05&     7.25$\\pm$0.05&  \\ldots & II  & X-ray, {\\em $K_s$}\\\\ MIR-64 & 21:40:44.83 &  58:15:03.3 &    14.62$\\pm$0.04  &    13.35$\\pm$ 0.04  &    12.66$\\pm$ 0.03 &     5.64$\\pm$0.02&     5.23$\\pm$0.02& 7.28$\\pm$0.05&    12.94$\\pm$0.06&  \\ldots & I/II & \\ha, X-ray\\\\ MIR-65 & 21:40:44.89 &  58:17:47.7 &   \\ldots  &     \\ldots  & \\ldots &     0.42$\\pm$0.00&     2.56$\\pm$0.01&     2.01$\\pm$0.02& 2.29$\\pm$0.03&  \\ldots & H$_2$\\\\ MIR-66\\tablenotemark{b} & 21:40:45.11 &  58:15:24.0 &   18.46$\\pm$0.02 &  16.74$\\pm$0.01 & 15.71$\\pm$0.01 &     0.59$\\pm$0.01&     0.44$\\pm$0.01&     1.73$\\pm$0.03& 5.26$\\pm$0.05&  \\ldots & I/II & \\\\ MIR-67 & 21:40:45.16 &  58:15:59.6 &    \\ldots  &     \\ldots  & \\ldots &     5.73$\\pm$0.02&     9.26$\\pm$0.03&    11.46$\\pm$0.06& 11.78$\\pm$0.07&  \\ldots & I& X-ray \\\\ MIR-68\\tablenotemark{b} & 21:40:45.24 &  58:16:41.3 &    17.38$\\pm$0.01 &  16.29$\\pm$ 0.01 &  15.53$\\pm$0.01 &     0.43$\\pm$0.01&     0.42$\\pm$0.01&     0.41$\\pm$0.02& 0.68$\\pm$0.02&  \\ldots & Galaxy\\\\ MIR-69 & 21:40:45.29 &  58:16:53.3 &   \\ldots  &     \\ldots  &     \\ldots &     0.53$\\pm$0.01&     0.61$\\pm$0.01&     0.60$\\pm$0.02& \\ldots&  \\ldots & III/S\\\\ MIR-70\\tablenotemark{b} & 21:40:45.41 &  58:16:45.5 &    20.38$\\pm$0.15 &   17.48$\\pm$0.02 &  15.96$\\pm$0.01 &  0.33$\\pm$0.00& 0.27$\\pm$0.01&     0.48$\\pm$0.02& 1.14$\\pm$0.02&  \\ldots & Galaxy\\\\ MIR-71 & 21:40:45.48 &  58:15:11.5 &    14.65$\\pm$0.05  &    13.71$\\pm$ 0.06  &    13.11$\\pm$ 0.04 &     3.17$\\pm$0.02&     2.61$\\pm$0.01& 3.41$\\pm$0.04&     5.73$\\pm$0.05&  \\ldots & I/II & X-ray, {\\em $K_s$}\\\\ MIR-72 & 21:40:45.50 &  58:16:02.6 &    15.68$\\pm$0.09  &    13.73$\\pm$ 0.06  &    12.85$\\pm$ 0.04 &     4.13$\\pm$0.02&     3.15$\\pm$0.02& 2.74$\\pm$0.03&     2.88$\\pm$0.03&  \\ldots & II & X-ray\\\\ MIR-73 & 21:40:45.76 &  58:15:48.9 &  \\ldots  &     \\ldots  & \\ldots &     2.71$\\pm$0.02&     7.27$\\pm$0.03&    10.79$\\pm$0.06& 11.29$\\pm$0.07&  \\ldots & I & X-ray\\\\ MIR-74 & 21:40:45.82 &  58:17:43.0 &   \\ldots  &     \\ldots  & \\ldots &     0.36$\\pm$0.00&     0.92$\\pm$0.01&     1.16$\\pm$0.02& 1.38$\\pm$0.02&  \\ldots & H$_2$\\\\ MIR-75 & 21:40:46.22 &  58:17:40.2 &  \\ldots  &     \\ldots  &     \\ldots &     0.34$\\pm$0.00&     0.71$\\pm$0.01&     1.20$\\pm$0.02& 1.53$\\pm$0.02&  \\ldots & H$_2$\\\\ MIR-76 & 21:40:46.46 &  58:15:23.0 &    12.81$\\pm$0.03  &    11.95$\\pm$ 0.03  &    11.65$\\pm$ 0.02 &     8.66$\\pm$0.03&     5.35$\\pm$0.02& 5.25$\\pm$0.04&     7.74$\\pm$0.07&  \\ldots & Trans. & X-ray, \\ha\\\\ MIR-77 & 21:40:46.86 &  58:15:33.4 &  15.30$\\pm$0.06  &    13.55$\\pm$ 0.05  &    12.63$\\pm$ 0.04 &     5.80$\\pm$0.02&     4.05$\\pm$0.02& 5.16$\\pm$0.04&    10.41$\\pm$0.08&  \\ldots & Galaxy & X-ray \\\\ MIR-78 & 21:40:47.35 &  58:16:37.3 &    17.08$\\pm$\\ldots & 16.14$\\pm$0.24 & 14.63$\\pm$0.10 &     1.00$\\pm$0.01&     0.75$\\pm$0.01& 0.59$\\pm$0.02&     0.36$\\pm$0.02&  \\ldots & III/S\\\\ MIR-79 & 21:40:47.43 &  58:17:39.1 &    11.53$\\pm$0.03  &     9.98$\\pm$ 0.03  &     9.37$\\pm$ 0.02 &    72.47$\\pm$0.17&    46.36$\\pm$0.10& 34.42$\\pm$0.11&    20.21$\\pm$0.07&  \\ldots & III/S\\\\ MIR-80 & 21:40:48.00 &  58:15:37.6 &     13.89$\\pm$0.03  &    12.95$\\pm$ 0.03  &    12.67$\\pm$ 0.03 &     4.72$\\pm$0.02&     3.39$\\pm$0.02& 5.77$\\pm$0.04&    11.90$\\pm$0.06&  \\ldots & I/II & \\ha, X-ray\\\\ MIR-81 & 21:40:48.52 &  58:16:37.1 &   16.18$\\pm$0.09  &    14.95$\\pm$ 0.08  &    14.46$\\pm$ 0.08 &     0.71$\\pm$0.01&     0.46$\\pm$0.01& 0.38$\\pm$0.02&     0.62$\\pm$0.02&  \\ldots & Galaxy& {\\em $K_s$} \\\\ MIR-82 & 21:40:48.58 &  58:16:00.9 &    15.88$\\pm$0.09  &    13.45$\\pm$ 0.04  &    11.88$\\pm$ 0.02 &    30.35$\\pm$0.08&    37.01$\\pm$0.08& 43.12$\\pm$0.13&    61.94$\\pm$0.17&    167.10$\\pm$0.09 & II& {\\em JH}\\\\ MIR-83 & 21:40:48.84 &  58:16:22.8 &   16.02$\\pm$0.12  &    15.25$\\pm$ 0.13  &    14.46$\\pm$ 0.08 &     0.61$\\pm$0.01&     0.43$\\pm$0.01& 0.25$\\pm$0.02& \\ldots&  \\ldots & III/S \\\\ MIR-84 & 21:40:48.87 &  58:15:10.9 &   15.64$\\pm$0.06  &    15.04$\\pm$ 0.08  &    14.56$\\pm$ 0.09 &     0.77$\\pm$0.01&     0.62$\\pm$0.01& 0.75$\\pm$0.02&     1.34$\\pm$0.03&  \\ldots & I/II & \\ha, {\\em $JK_s$}\\\\ MIR-85 & 21:40:49.05 &  58:17:19.4 &   15.79$\\pm$0.08  &    13.83$\\pm$ 0.05  &    12.92$\\pm$ 0.02 &     3.32$\\pm$0.02&     2.19$\\pm$0.01& 1.75$\\pm$0.02&     1.18$\\pm$0.02&  \\ldots & Trans. & {\\em J} \\\\ MIR-86 & 21:40:49.06 &  58:17:09.3 &     14.14$\\pm$0.03  &    12.86$\\pm$ 0.04  &    12.13$\\pm$ 0.02 &    13.88$\\pm$0.05&    14.57$\\pm$0.04& 15.54$\\pm$0.07&    17.83$\\pm$0.07&     17.99$\\pm$0.05 & II & \\ha, X-ray, {\\em $K_s$}\\\\ MIR-87 & 21:40:49.29 &  58:14:11.3 &  13.52$\\pm$0.03  &    12.66$\\pm$ 0.03  &    12.35$\\pm$ 0.02 &     3.87$\\pm$0.02&     2.44$\\pm$0.01& 1.66$\\pm$0.02&     0.99$\\pm$0.02&  \\ldots & III/S\\\\ MIR-88 & 21:40:49.67 &  58:17:37.3 &   17.13$\\pm$0.23  &    14.99$\\pm$ 0.09  &    14.23$\\pm$ 0.07 &     1.07$\\pm$0.01&     0.73$\\pm$0.01& 0.60$\\pm$0.02&     0.36$\\pm$0.02&  \\ldots & III/S \\\\ MIR-89\\tablenotemark{b} & 21:40:51.07 &  58:16:24.3 & 17.80$\\pm$0.01 &  16.70$\\pm$0.01 & 16.09$\\pm$0.01 &   0.11$\\pm$0.00& 0.08$\\pm$0.00&     0.11$\\pm$0.02& 0.23$\\pm$0.02&  \\ldots & Galaxy \\\\ MIR-90 & 21:40:51.23 &  58:16:38.9 &    14.23$\\pm$0.03  &    13.76$\\pm$ 0.04  &    13.61$\\pm$ 0.04 &     1.26$\\pm$0.01&     0.78$\\pm$0.01& 0.57$\\pm$0.02&     0.53$\\pm$0.02&  \\ldots & Trans. &{\\em HK} \\\\ MIR-91 & 21:40:51.81 &  58:16:29.1 &   15.70$\\pm$0.06  &    14.45$\\pm$ 0.06  &    13.93$\\pm$ 0.07 &     1.19$\\pm$0.01&     0.74$\\pm$0.01& 0.60$\\pm$0.02&     0.54$\\pm$0.02&  \\ldots & Trans. &{\\em HK} \\\\ MIR-92 & 21:40:51.90 &  58:17:01.6 &   \\ldots  &     \\ldots  &     \\ldots &     0.36$\\pm$0.00&     0.33$\\pm$0.00&     0.20$\\pm$0.02& \\ldots&  \\ldots & III/S\\\\ MIR-93\\tablenotemark{b} & 21:40:51.97 &  58:16:23.6 &  19.21$\\pm$0.06 & 17.47$\\pm$0.01 & 16.39$\\pm$0.01 & 0.16$\\pm$0.00& 0.10$\\pm$0.00&     0.23$\\pm$0.02& 0.65$\\pm$0.02&  \\ldots &Galaxy\\\\ MIR-94 & 21:40:52.26 &  58:17:46.0 &    15.68$\\pm$0.05  &    14.56$\\pm$ 0.05  &    14.06$\\pm$ 0.06 &     0.94$\\pm$0.01&     0.65$\\pm$0.01& 0.59$\\pm$0.02&     0.48$\\pm$0.03&  \\ldots & Trans. \\\\ MIR-95 & 21:40:52.98 &  58:18:01.1 &    15.29$\\pm$0.05  &    14.59$\\pm$ 0.06  &    14.22$\\pm$ 0.07 &     0.84$\\pm$0.01&     0.53$\\pm$0.01& 0.54$\\pm$0.03&     0.34$\\pm$0.03&  \\ldots &III/S \\\\ MIR-96 & 21:40:53.12 &  58:14:57.4 & 14.77$\\pm$0.06  &    14.25$\\pm$ 0.07  &    13.95$\\pm$ 0.07 &     0.84$\\pm$0.01&     0.54$\\pm$0.01& 0.38$\\pm$0.02&     0.37$\\pm$0.03&  \\ldots & Trans. \\\\ MIR-97 & 21:40:53.42 &  58:14:51.9 &     15.50$\\pm$0.06  &    14.88$\\pm$ 0.09  &    14.74$\\pm$ 0.11 &     0.44$\\pm$0.00&     0.31$\\pm$0.01& 0.23$\\pm$0.01&     0.20$\\pm$0.02&  \\ldots & Trans.\\\\ MIR-98 & 21:40:53.68 &  58:16:41.2 &   \\ldots  &     \\ldots  &     \\ldots &     0.25$\\pm$0.00&     0.71$\\pm$0.01&     0.79$\\pm$0.02& 0.83$\\pm$0.02&  \\ldots & H$_2$ \\\\ MIR-99 & 21:40:55.68 &  58:15:50.8 &   \\ldots  &     \\ldots  & \\ldots &     0.36$\\pm$0.00&     0.14$\\pm$0.00&     2.69$\\pm$0.03& 9.65$\\pm$0.05&  \\ldots & I/II \\\\ MIR-100 & 21:40:55.87 &  58:14:16.3 &     14.18$\\pm$0.03  &    13.61$\\pm$ 0.04  &    13.39$\\pm$ 0.04 & \\ldots&     0.87$\\pm$0.02& \\ldots& 0.45$\\pm$0.03&  \\ldots & III/S \\\\ MIR-101 & 21:40:55.88 &  58:15:34.8 &   \\ldots  &     \\ldots  & \\ldots &     0.25$\\pm$0.00&     0.32$\\pm$0.00&     0.41$\\pm$0.02& 0.71$\\pm$0.02&  \\ldots & Galaxy\\\\ MIR-102 & 21:40:55.93 &  58:17:07.8 &  14.45$\\pm$0.03  &    13.19$\\pm$ 0.04  &    12.76$\\pm$ 0.03 &     3.21$\\pm$0.02&     2.01$\\pm$0.01& 1.39$\\pm$0.02&     0.99$\\pm$0.02&  \\ldots & Trans. \\\\ MIR-103 & 21:40:57.04 &  58:16:56.8 &    16.61$\\pm$0.13  &    15.92$\\pm$ 0.21  &    16.22$\\pm$ 0.00 &     0.42$\\pm$0.01&     0.34$\\pm$0.00& 0.40$\\pm$0.02&     1.19$\\pm$0.03&  \\ldots & Galaxy \\\\ MIR-104 & 21:40:57.06 &  58:16:29.0 &  14.73$\\pm$0.03  &    13.00$\\pm$ 0.03  &    12.24$\\pm$ 0.02 &     6.10$\\pm$0.02&     4.11$\\pm$0.02& 2.87$\\pm$0.03&     1.96$\\pm$0.03&  \\ldots & Trans. \\\\ MIR-105 & 21:40:57.33 &  58:16:41.7 &   16.64$\\pm$\\ldots & 15.41$\\pm$\\ldots  & 14.55$\\pm$ 0.13  &    0.79$\\pm$0.01& 1.18$\\pm$0.01& 2.90$\\pm$0.03&     6.34$\\pm$0.04&  \\ldots & H$_2$\\\\ MIR-106 & 21:40:57.50 &  58:14:43.6 &     15.00$\\pm$0.04  &    14.66$\\pm$ 0.07  &    14.34$\\pm$ 0.09 & \\ldots&     0.43$\\pm$0.01& 0.25$\\pm$0.03&     0.15$\\pm$0.02&  \\ldots & III/S\\\\ MIR-107 & 21:40:58.03 &  58:15:12.1 &    15.06$\\pm$0.05  &    14.50$\\pm$ 0.07  &    14.12$\\pm$ 0.08 &     0.70$\\pm$0.01&     0.46$\\pm$0.01& 0.30$\\pm$0.02&     0.21$\\pm$0.02&  \\ldots & Trans. \\\\ MIR-108 & 21:40:58.45 &  58:16:20.5 &     13.11$\\pm$0.03  &    12.44$\\pm$ 0.03  &    12.02$\\pm$ 0.03 &     5.93$\\pm$0.03&     3.77$\\pm$0.02& 3.13$\\pm$0.03&     2.52$\\pm$0.03&  \\ldots & Trans. \\\\ MIR-109 & 21:40:58.53 &  58:15:00.3  & 12.15$\\pm$0.03  &    11.26$\\pm$ 0.03  &    10.96$\\pm$ 0.02 & \\ldots&     8.34$\\pm$0.03& 5.73$\\pm$0.05&     3.46$\\pm$0.03&  \\ldots  & III/S\\\\ MIR-110 & 21:40:59.71 &  58:17:33.3  &   12.69$\\pm$0.02  &    11.25$\\pm$ 0.03  &    10.59$\\pm$ 0.02 &    24.86$\\pm$0.17& \\ldots& 10.77$\\pm$0.13& \\ldots&  \\ldots & III/S \\\\ \\enddata \\tablenotetext{a}{$JHK_s$ photometry from \\citet{nisini01}} \\tablenotetext{b}{$JHK_s$ photometry from \\citet{beltran09}} \\tablenotetext{c}{Variable in $J/H/K_s$ } \\end{deluxetable}  \\setcounter{table}{1} \\begin{deluxetable}{lccccccc} \\tabletypesize{\\scriptsize} \\tablecaption{Coordinates and flux densities of NIR \\& IRAC 3.6/4.5 $\\mu$m PMS sources in SFO~38 \\label{tab_nirpms}} \\tablewidth{0pt} \\tablehead{ \\colhead{Source} & \\colhead{$\\alpha_{2000}$}& \\colhead{$\\delta_{2000}$}& \\colhead{$J$}& \\colhead{$H$}& \\colhead{$K_{\\rm s}$}& \\colhead{$F_{3.6}$}& \\colhead{$F_{4.5}$}\\\\ SFO38 & & & &  & &mJy &mJy \\\\} \\startdata NIR-1 & 21:40:34.36 &  58:16:19.2 & 20.88$\\pm$\\ldots & 18.12$\\pm$0.03 & 16.68$\\pm$0.02 & 15.73$\\pm$0.03 & 15.38$\\pm$0.03\\\\ NIR-2 & 21:40:36.46 &  58:16:28.7 & 20.84$\\pm$\\ldots & 19.31$\\pm$0.10 & 17.14$\\pm$0.03 & 15.23$\\pm$0.02 & 14.75$\\pm$0.02\\\\ NIR-3 & 21:40:37.57 &  58:16:01.2 & 20.80$\\pm$\\ldots & 19.98$\\pm$\\ldots & 17.35$\\pm$\\ldots & 15.22$\\pm$0.03 & 14.29$\\pm$0.02\\\\ NIR-4 & 21:40:38.17 &  58:15:32.7 & 20.87$\\pm$\\ldots & 18.32$\\pm$0.03 & 16.72$\\pm$0.02 & 14.84$\\pm$0.02 & 14.59$\\pm$0.02\\\\ NIR-5 & 21:40:39.53 &  58:15:46.9 &  \\ldots &  \\ldots &   \\ldots & 14.22$\\pm$0.01 & 13.72$\\pm$0.01\\\\ NIR-6 & 21:40:39.65 &  58:15:06.6 & 16.59$\\pm$0.01 & 15.88$\\pm$0.01 & 15.52$\\pm$0.01 & 14.07$\\pm$0.01 & 13.82$\\pm$0.01\\\\ NIR-7 & 21:40:42.81 &  58:15:45.3 & 16.29$\\pm$0.00 & 15.38$\\pm$0.00 & 14.77$\\pm$0.00 & 13.91$\\pm$0.01 & 13.27$\\pm$0.01\\\\ NIR-8 & 21:40:44.15 &  58:15:49.6 &  \\ldots &  \\ldots &   \\ldots & 15.78$\\pm$0.04 & 13.83$\\pm$0.01\\\\ NIR-9 & 21:40:44.57 &  58:15:51.6 & 15.97$\\pm$0.00 & 15.23$\\pm$0.00 & 14.71$\\pm$0.00 & 14.28$\\pm$0.02 & 13.68$\\pm$0.01\\\\ NIR-10 & 21:40:45.85 &  58:15:40.6 & 18.70$\\pm$0.03 & 17.09$\\pm$0.01 & 16.01$\\pm$0.01 & 14.68$\\pm$0.02 & 14.18$\\pm$0.02\\\\ NIR-11 & 21:40:48.28 &  58:16:18.8 & 18.17$\\pm$0.02 & 17.16$\\pm$0.01 & 16.40$\\pm$0.02 & 15.63$\\pm$0.03 & 15.24$\\pm$0.03\\\\ NIR-12 & 21:40:50.55 &  58:16:01.8 & 20.84$\\pm$\\ldots & 19.65$\\pm$0.10 & 17.55$\\pm$0.03 & 15.65$\\pm$0.03 & 15.29$\\pm$0.03\\\\ NIR-13 & 21:40:51.14 &  58:17:01.5 & 20.86$\\pm$\\ldots & 18.74$\\pm$0.05 & 17.23$\\pm$0.04 & 15.99$\\pm$0.03 & 15.68$\\pm$0.04\\\\ \\enddata \\end{deluxetable}  \\setcounter{table}{2} \\begin{deluxetable}{lccccc} \\tabletypesize{\\scriptsize} \\tablecaption{{\\em BVRI} photometry of MIR sources in SFO~38 \\label{tab_optmag}} \\tablewidth{0pt} \\tablehead{ \\colhead{Source} & \\colhead{$B$}& \\colhead{$V$}& \\colhead{$R$}& \\colhead{$I$}& \\colhead{Variability}\\\\} \\startdata MIR-1 &  19.82$\\pm$0.010  & 18.24$\\pm$0.005  &17.18$\\pm$0.004  &16.00$\\pm$0.003  &        \\\\ MIR-2 &  20.48$\\pm$0.015  & 18.88$\\pm$0.006  &17.88$\\pm$0.008  &16.84$\\pm$0.005  &        \\\\ MIR-3 &  21.58$\\pm$0.021  & 19.58$\\pm$0.007  &18.37$\\pm$0.006  &17.14$\\pm$0.004  &        \\\\ MIR-4 &   0.00$\\pm$0.000  & 20.38$\\pm$0.016  &19.15$\\pm$0.007  &17.73$\\pm$0.006  &        \\\\ MIR-5 &  21.11$\\pm$0.015  & 19.32$\\pm$0.007  &17.77$\\pm$0.004  &16.37$\\pm$0.002  & {\\em BVI}       \\\\ MIR-7 &  21.59$\\pm$0.021  & 19.39$\\pm$0.006  &17.90$\\pm$0.004  &16.38$\\pm$0.002  &        \\\\ MIR-8 &  17.74$\\pm$0.001  & 15.74$\\pm$0.001  &14.55$\\pm$0.001  &13.35$\\pm$0.001  &        \\\\ MIR-9 &  19.95$\\pm$0.017  & 18.48$\\pm$0.004  &17.60$\\pm$0.005  &16.62$\\pm$0.004  &        \\\\ MIR-10 &  17.28$\\pm$0.001  & 16.07$\\pm$0.001  &15.29$\\pm$0.001  &14.49$\\pm$0.001  &        \\\\ MIR-11 &   0.00$\\pm$0.000  & 19.52$\\pm$0.011  &17.94$\\pm$0.007  &16.19$\\pm$0.005  &        \\\\ MIR-14 &  17.67$\\pm$0.001  & 16.24$\\pm$0.001  &15.38$\\pm$0.001  &14.53$\\pm$0.001  &        \\\\ MIR-15 &  16.89$\\pm$0.001  & 15.26$\\pm$0.001  &14.33$\\pm$0.001  &13.44$\\pm$0.001  &        \\\\ MIR-18 &  19.62$\\pm$0.006  & 17.58$\\pm$0.003  &16.30$\\pm$0.002  &14.98$\\pm$0.009  &        \\\\ MIR-26 &  17.34$\\pm$0.001  & 16.03$\\pm$0.001  &15.20$\\pm$0.001  &14.30$\\pm$0.001  &        \\\\ MIR-29 &  20.58$\\pm$0.011  & 18.43$\\pm$0.005  &16.94$\\pm$0.002  &15.42$\\pm$0.002  &        \\\\ MIR-31 &  17.43$\\pm$0.002  & 15.71$\\pm$0.001  &14.60$\\pm$0.001  &13.52$\\pm$0.001  &        \\\\ MIR-32 &  21.15$\\pm$0.015  & 19.29$\\pm$0.006  &18.01$\\pm$0.007  &16.25$\\pm$0.002  & {\\em BVI}\\\\ MIR-43 &  19.41$\\pm$0.004  & 17.49$\\pm$0.002  &16.21$\\pm$0.001  &14.73$\\pm$0.001  & {\\em BVI}\\\\ MIR-44 &  19.36$\\pm$0.006  & 17.65$\\pm$0.003  &16.57$\\pm$0.001  &15.51$\\pm$0.002  &        \\\\ MIR-45 &  21.41$\\pm$0.018  & 19.45$\\pm$0.007  &18.10$\\pm$0.012  &16.33$\\pm$0.002  & {\\em BVI}\\\\ MIR-47 &   0.00$\\pm$0.000  & 19.86$\\pm$0.009  &18.76$\\pm$0.005  &17.61$\\pm$0.005  &        \\\\ MIR-49 &  20.91$\\pm$0.012  & 18.81$\\pm$0.004  &17.04$\\pm$0.003  &15.69$\\pm$0.001  & {\\em BVI}\\\\ MIR-52 &  19.54$\\pm$0.008  & 18.32$\\pm$0.005  &17.52$\\pm$0.003  &16.64$\\pm$0.005  &        \\\\ MIR-58 &  20.13$\\pm$0.007  & 18.75$\\pm$0.003  &17.86$\\pm$0.005  &16.91$\\pm$0.003  & {\\em B}  \\\\ MIR-60 &  20.79$\\pm$0.011  & 19.13$\\pm$0.005  &18.07$\\pm$0.005  &17.06$\\pm$0.003  &        \\\\ MIR-64 &   0.00$\\pm$0.000  & 20.49$\\pm$0.014  &18.77$\\pm$0.016  &17.09$\\pm$0.006  &        \\\\ MIR-76 &  19.08$\\pm$0.006  & 17.16$\\pm$0.002  &15.90$\\pm$0.002  &14.55$\\pm$0.001  &        \\\\ MIR-79 &  21.41$\\pm$0.024  & 18.56$\\pm$0.004  &16.52$\\pm$0.002  &14.55$\\pm$0.001  &        \\\\ MIR-80 &  21.69$\\pm$0.023  & 19.81$\\pm$0.008  &17.90$\\pm$0.006  &16.00$\\pm$0.004  &{\\em V}   \\\\ MIR-87 &  19.72$\\pm$0.005  & 17.72$\\pm$0.002  &16.50$\\pm$0.002  &15.28$\\pm$0.002  &        \\\\ MIR-90 &  18.15$\\pm$0.002  & 16.85$\\pm$0.002  &16.04$\\pm$0.001  &15.26$\\pm$0.002  &        \\\\ MIR-95 &   0.00$\\pm$0.000  & 20.16$\\pm$0.012  &18.71$\\pm$0.009  &17.34$\\pm$0.004  &        \\\\ MIR-96 &  19.23$\\pm$0.008  & 17.85$\\pm$0.003  &16.94$\\pm$0.003  &16.02$\\pm$0.002  &        \\\\ MIR-97 &  20.06$\\pm$0.007  & 18.63$\\pm$0.005  &17.76$\\pm$0.005  &16.83$\\pm$0.004  &        \\\\ MIR-100 &  18.31$\\pm$0.003  & 16.86$\\pm$0.002  &16.03$\\pm$0.001  &15.28$\\pm$0.002  &        \\\\ MIR-106 &  18.87$\\pm$0.004  & 17.75$\\pm$0.003  &16.96$\\pm$0.003  &16.21$\\pm$0.002  &        \\\\ MIR-107 &  19.50$\\pm$0.004  & 18.12$\\pm$0.002  &17.22$\\pm$0.003  &16.27$\\pm$0.002  &        \\\\ MIR-108 &  18.94$\\pm$0.005  & 17.31$\\pm$0.004  &16.13$\\pm$0.003  &14.86$\\pm$0.001  &        \\\\ MIR-109 &  18.17$\\pm$0.003  & 16.17$\\pm$0.001  &15.00$\\pm$0.001  &13.86$\\pm$0.001  &        \\\\ MIR-110 &   0.00$\\pm$0.000  & 19.78$\\pm$0.010  &17.76$\\pm$0.004  &15.76$\\pm$0.003  &{\\em V}   \\\\ BRC38 1 &   0.00$\\pm$0.000  & 20.05$\\pm$0.017  &18.24$\\pm$0.009  &16.45$\\pm$0.003  &        \\\\ BRC38 17 &   0.00$\\pm$0.000  & 20.51$\\pm$0.017  &18.65$\\pm$0.005  &16.92$\\pm$0.006  & {\\em V} \\\\ GFG 64 &  19.01$\\pm$0.007  & 17.48$\\pm$0.003  &16.19$\\pm$0.001  &15.18$\\pm$0.001  &        \\\\ IPHAS &   0.00$\\pm$0.000  & 20.28$\\pm$0.011  &18.28$\\pm$0.009  &16.50$\\pm$0.004  &        \\\\ \\enddata \\end{deluxetable}  \\setcounter{table}{3} \\begin{deluxetable}{llllllllllllc} \\tabletypesize{\\scriptsize} \\rotate \\tablecaption{Spectral Classification of YSOs \\label{tab_sptclass}} \\tablewidth{0pt} \\tablehead{ \\colhead{ID}& \\colhead{Alt. ID} & \\colhead{Date}& \\colhead{W$_\\lambda$}& \\colhead{Sp. Type}& \\colhead{T$_{eff}$} & \\colhead{Type}& \\colhead{\\ha[10$\\%$]}& \\colhead{log \\.{M}$_{ac}$}& \\colhead{$\\alpha_{IRAC}$} & \\colhead{E(V-I)} & \\colhead{A$_{V}$}& \\colhead{E(V-I)/E(B-V)}\\\\ &           &          & (\\AA)     &($\\pm$subtype)& (K)&     & (km s$^{-1})$  &(\\msun yr$^{-1}$)    &          &       &        &      \\\\} \\startdata BRC 38 1&           &   5/10/08&  -12.93 & M3$\\pm$1&  3470&  CTTS&  312.23 &  -9.86   &         &  1.13 &   2.81 &       \\\\ &           &  10/10/09&  -16.24 &         &      &      &  205.29 &  -10.90  &         &       &        &       \\\\ MIR-5   &  BRC 38 2 &   5/10/08&  -85.23 & M1$\\pm$1&  3720&  CTTS&  359.79 &  -9.40   &         &  0.99 &   2.47 &  3.19 \\\\ &           &  10/10/09&  -106.1 &         &      &      &  244.53 &  -10.52  &         &       &        &       \\\\ MIR-7   &  BRC 38 3 &   6/10/08&  -19.0  & M1$\\pm$2&  3720&  CTTS&  476.25 &  -8.27   & -0.87   &  1.05 &   2.61 &  1.46 \\\\ &           &  10/10/09&  -18.75 &         &      &      &  463.50 &  -8.39   &         &       &        &       \\\\ MIR-11  &  BRC 38 4 &  10/09/08&  -7.30  & M1$\\pm$2&  3720&  WTTS&         &          &  0.02   &  1.37 &   3.41 &       \\\\ &           &  18/07/09&  -11.29 &         &      &      &         &          &         &       &        &       \\\\ MIR-29  &  BRC 38 5 &   6/09/08&  -3.42  & M1$\\pm$2&  3720&  WTTS&         &          & -1.74   &  1.05 &   2.61 &  1.57 \\\\ &           &   3/11/08&  -8.84  &         &      &\t &  469.90 &  -8.33   &         &       &        &       \\\\ &           &  24/08/09&  -1.89  &         &      &      & \t   &          &         &       &        &       \\\\ &           &   9/11/09&  -2.74  &         &      &      &         &          &         &       &        &       \\\\ MIR-31  &  BRC 38 6 &   6/09/08&  -53.47 & K2$\\pm$1&  4900&  CTTS&  384.43 &  -9.16   & -1.60   &  1.18 &   2.94 &  1.42 \\\\ &           &  24/08/09&  -57.32 &         &      &      &  442.02 &  -8.60   &         &       &        &       \\\\ &           &   9/10/09&  -60.5  &         &      &      &  382.26 &  -9.18   &         &       &        &       \\\\ MIR-32  &  BRC 38 7 &  10/09/08&  -33.12 & M2$\\pm$1&  3580&  CTTS&  368.03 &  -9.32   & -1.03   &  0.90 &   2.24 &  2.65 \\\\ &\t    & 12/09/09&  -15.78  &         &      &      &  220.59 &  -10.75  &         & \t&  \t &       \\\\ MIR-43  &  GFG 62   &  18/07/09&  -3.516 & M1$\\pm$1&  3720&  WTTS&         &          & -2.61   &  0.80 &   1.99 &  1.82 \\\\ MIR-45  &  BRC 38 9 &   6/09/08&  -37.43 & K4$\\pm$1&  4590&  CTTS&  281.09 &  -10.16  & -0.18   &  1.97 &   4.91 &  2.21 \\\\ &           &   3/11/08&  -41.22 &         &      &      &  389.74 &  -9.11   &         &       &        &       \\\\ &           &  24/08/09&  -56.33 &         &      &      &  273.39 &  -10.24  &         & \t&  \t &       \\\\ MIR-49  &  BRC 38 10&   4/11/08&  -13.5  & M0$\\pm$1&  3850&  CTTS&  430.70 &  -8.71   & -1.37   &  1.32 &   3.29 &  1.91 \\\\ &\t    &  9/10/09&  -30     &         &      &      &  450.55 &  -8.52   & \t& \t&  \t &       \\\\ MIR-64\t&  BRC 38 11&  10/09/08&  -31.46 & K5$\\pm$1&  4350&  CTTS&  239.20 &  -10.57  &  0.10   &  2.04 &   5.08 &       \\\\ &\t    & 12/09/09&  -25.06  &         &      &      &         &          &  \t&       &        &       \\\\ MIR-76  &  GFG 81   &  18/07/09&  -3.261 & K7$\\pm$1&  4060&  WTTS&         &          & -1.07   &  1.01 &   2.51 &  1.87 \\\\ &\t    & 10/10/09&  -3.135  &         &      &      &         &          &  \t&       &        &       \\\\ MIR-80  &  BRC 38 12&  24/08/09&  -8.404 & M4$\\pm$1&  3370&  WTTS&         &          &  0.31   &  0.95 &   2.37 &  3.39 \\\\ BRC 38 17&          &  11/11/09&  -119.  & M3$\\pm$2&  3470&  CTTS&  205.29 &  -10.90  &         &  1.12 &   2.79 &       \\\\ \\enddata \\end{deluxetable} \\setcounter{table}{4} \\begin{deluxetable}{cccccccccc} \\tabletypesize{\\small} \\tablecolumns{9} \\tablewidth{0pt} \\tablecaption{Parameters derived from SED modeling using axisymmetric radiation transfer models for the candidate YSOs.  \\av\\ refers to the foreground extinction towards the source. \\label{tab_sedfits}} \\tablehead{ \\colhead{Source} & \\colhead{$T_{\\ast}$} & \\colhead{\\mstar} & \\colhead{$M_{\\rm disk}$} & \\colhead{$M_{\\rm env}$} & \\colhead{$\\dot{M}_{\\rm env}$} & \\colhead{$\\dot{M}_{\\rm disk}$} & \\colhead{$L$} & \\colhead{\\av} & \\colhead{Age}\\\\ \\colhead{}& \\colhead{K} & \\colhead{\\msun} & \\colhead{\\rm M$_{\\odot}$} & \\colhead{\\rm M$_{\\odot}$} & \\colhead{\\rm M$_{\\odot}$~yr$^{-1}$} &\\colhead{\\rm M$_{\\odot}$~yr$^{-1}$} & \\colhead{\\rm L$_{\\odot}$} & & Myr } \\startdata MIR-11&  4442 &  1.26  & 6.44$\\times10^{-3}$ &  6.39$\\times10^{-5}$  & 0.0                 & 2.10$\\times10^{-9}$ &1.27 &  5.3 & 3.2\\\\ MIR-29&  4837 &  1.55  & 1.66$\\times10^{-3}$ &  6.13$\\times10^{-7}$  & 0.0                 & 1.16$\\times10^{-9}$ &1.67   &  4.4 & 5.8\\\\ MIR-31&  5923 &  1.88  & 1.33$\\times10^{-4}$ &  7.04$\\times10^{-8}$  & 0.0                 & 2.41$\\times10^{-9}$ &8.92   &  4.2 & 6.7\\\\ MIR-34&  4523 &  1.53  & 7.64$\\times10^{-5}$ &  2.56$\\times10^{-3}$  & 1.21$\\times10^{-8}$ & 1.28$\\times10^{-10}$& 5.8 &  3.2  &0.69\\\\ MIR-36&  5413 &  3.29  & 4.37$\\times10^{-3}$ &  1.85$\\times10^{-8}$  & 0.0                 & 4.80$\\times10^{-8}$ &29.1 &  49.6 & 1.3\\\\ MIR-38&  7854 &  1.98  & 3.10$\\times10^{-5}$ &  2.99$\\times10^{-6}$  & 0.0                 & 5.49$\\times10^{-11}$&11.9  &  23.0 & 8.2\\\\ MIR-45&  9960 &  2.32  & 4.23$\\times10^{-5}$ &  1.70$\\times10^{-8}$  & 0.0                 & 1.37$\\times10^{-9}$ &27.0 &  8.3 & 9.4\\\\ MIR-48&  7554 &  5.14  & 1.71$\\times10^{-3}$ &  21.4$\\times10^{-2}$  & 2.23$\\times10^{-7}$ & 2.38$\\times10^{-9}$ & 280 &45.9 & 0.49\\\\ MIR-50&  4643 &  5.97  & 7.67$\\times10^{-2}$ &  40.8                 & 3.82$\\times10^{-4}$ & 2.08$\\times10^{-8}$ &196.7   &  44.4 &0.1\\\\ MIR-54&  4290 &  1.50  & 2.46$\\times10^{-2}$ &  9.3                  & 7.13$\\times10^{-5}$ & 2.98$\\times10^{-6}$ &33.4  &  19.6 &0.17\\\\ MIR-55&  4989 &  3.07  & 7.86$\\times10^{-4}$ &  1.44$\\times10^{-3}$ & 1.97$\\times10^{-9}$ & 1.14$\\times10^{-9}$ &14.4   &  52.3 &0.91\\\\ MIR-59&  5218 &  4.27  & 9.99$\\times10^{-3}$ &  7.37$\\times10^{-3}$ & 2.95$\\times10^{-7}$ & 2.00$\\times10^{-8}$ & 60.5  &  87.4 & 0.53\\\\ MIR-82&  11320&  2.77  & 5.85$\\times10^{-5}$ &  1.18$\\times10^{-6}$  & 0.0                 & 1.40$\\times10^{-10}$ &57.2&  21.0 & 4.8\\\\ MIR-86&  4594 &  1.55  & 8.18$\\times10^{-4}$ &  3.63$\\times10^{-7}$  & 0.0                 & 4.50$\\times10^{-10}$ &2.59  &  9.7 & 1.8\\\\ \\enddata \\end{deluxetable}"
    },
    "1005/1005.3769_arXiv.txt": {
        "abstract": "We study the bulk viscosity of dense matter, taking into account non-linear effects which arise in the large amplitude ``supra-thermal'' region where the deviation $\\mu_\\Delta$ of the chemical potentials from chemical equilibrium fulfills $\\mu_\\Delta\\gtrsim T$. This regime is relevant to unstable modes such as r-modes, which grow in amplitude until saturated by non-linear effects. We study the damping due to direct and modified Urca processes in hadronic matter, and due to nonleptonic weak interactions in strange quark matter. We give general results valid for an arbitrary equation of state of dense matter and find that the viscosity can be strongly enhanced by supra-thermal effects. Our study confirms previous results on quark matter and shows that the non-linear enhancement is even stronger in the case of hadronic matter. Our results can be applied to calculations of the r-mode-induced spin-down of fast-rotating neutron stars, where the spin-down time will depend on the saturation amplitude of the r-mode. ",
        "introduction": "Compact stars are the only known objects that contain equilibrated matter that is compressed beyond nuclear density, making them a valuable laboratory for the study of the structure of matter under extreme conditions. In addition to hadronic matter they may also contain new forms of matter that involve deconfined quarks \\cite{Itoh:1970uw,Bodmer:1971we,Witten:1984rs,Alford:2007xm}. In contrast to the static properties of compact stars which only depend on the equation of state  of matter \\cite{Lattimer:2000nx}, dynamic properties also depend on the low energy degrees of freedom and thereby might be able to discriminate more efficiently between different forms of strongly interacting matter. One of the dynamic properties of dense matter is viscosity, which determines the damping of mechanical perturbations, and a particularly important application is to the damping of r-mode oscillations of compact stars \\cite{Papaloizou:1978,Andersson:1997xt,Friedman:1997uh,Lindblom:1999yk,Madsen:1999ci}, which, at sufficiently low viscosity and high rotation rate, are unstable and can cause rapid spin-down of the star \\cite{Owen:1998xg} via gravitational radiation. Since the r-mode is unstable, its exponential growth must eventually be stopped by some non-linear mechanism. Finding the relevant mechanism is important because it determines the amplitude at which the r-mode saturates, and hence the rate at which it spins down the star. Previously suggested mechanisms include mode coupling and the transformation of the r-mode energy into differential rotation \\cite{Lindblom:2000az,Gressman:2002zy,Lin:2004wx,Bondarescu:2007jw}, and friction between different layers of the star such as ``surface rubbing'' at the crust \\cite{Bildsten:2000}. Because of the complexity of the problem, these mechanisms have to rely on approximations that are not always well controlled.  In this paper we consider an alternative mechanism which does not involve additional physics, but is already present in a quasi-static hydrodynamic description. At low amplitudes the bulk viscosity is amplitude-independent, but since the r-mode is unstable its amplitude grows, and unless stopped by other mechanisms will quickly enter the ``supra-thermal'' regime where the bulk viscosity grows with amplitude, and may become large enough to stop the growth of the mode \\cite{Madsen:1992sx,Reisenegger:2003pd}. The supra-thermal regime is characterized by $\\mu_\\Delta\\gtrsim T$, where $T$ is the temperature and $\\mu_\\Delta$  is the amplitude of the oscillations in the chemical potential of the quantity whose re-equilibration causes the viscous damping. We will study the microscopic part of the problem via a comprehensive analysis of the bulk viscosity of dense matter. We leave the astrophysical aspects for future work. Since the precise phase structure and the equation of state of matter at high density is still unknown, we keep the dependence on those parameters as explicit as possible, and as well as numeric results we provide analytic approximations which prove to be surprisingly accurate. This allows us to obtain general results for the bulk viscosity valid for many different phases of matter, and enables us to estimate the involved uncertainties. We study in detail the cases of equilibrated $npe$-matter and strange quark matter, and consider both modified and direct Urca processes in the hadronic case. Yet, our general expressions can be applied to other equations of state and entirely different forms of strongly interacting matter. ",
        "conclusions": "We have studied the bulk viscosity of dense matter including its non-linear behavior at large amplitudes. In particular we give a general solution for the bulk viscosity of degenerate matter that is valid for arbitrary equations of state and retains its full parameter dependence in sec.~\\ref{sec:general-viscosity}. This allows one to include these supra-thermal effects in a systematic r-mode analysis. In the supra-thermal regime we give a general analytic result for strange quark matter with non-leptonic processes in sec.~\\ref{sub:Strange-quark-matter}. We found that the free hadron gas model of nuclear matter, used for example in \\cite{Lindblom:1998wf,Owen:1998xg,Jaikumar:2008kh} to compute the susceptibilities that enter the viscosity, is not accurate even in the sub-thermal regime, and may significantly underestimate the viscosity. Moreover, we find that the standard low temperature (high frequency) approximation is not applicable for temperatures around $10^{10}$~K and the full resonant form of the bulk viscosity is required. We confirm previous results for the amplitude-dependence of the bulk viscosity of strange quark matter \\cite{Madsen:1992sx} and find that these supra-thermal effects are parametrically even more important in nuclear matter, because of higher-order non-linearities in the amplitude-dependence of the Urca rate.  The most obvious application of our results is to the damping of unstable r-mode oscillations in neutron stars. As the amplitude of the mode enters the supra-thermal regime the viscosity will increase steeply above the sub-thermal result and can exceed it by many orders of magnitude, but eventually it reaches an upper bound that is completely independent of the particular weak damping process and depends only on susceptibilities of the dense matter in question. The viscosity then decreases at even larger amplitudes. We conclude that if r-mode growth is not stopped by the supra-thermal bulk viscosity before this maximum is reached then other non-linear dynamic effects \\cite{Lindblom:2000az,Gressman:2002zy,Lin:2004wx,Bondarescu:2007jw} will be required to stop it. We have already performed initial exploratory calculations of r-mode damping times, and these suggest that over a significant region of parameter space supra-thermal bulk viscosity is sufficient to saturate r-mode growth at a finite amplitude parameter $\\alpha_{max}<1$, as was previously assumed \\cite{Owen:1998xg}. This topic will be discussed in more detail in a forthcoming publication. There are several other directions in which our research could be developed. Other equations of state for quark matter could be studied, for example the perturbative equation of state \\cite{Kurkela:2009gj}, and also other phases with different equilibration mechanisms. The same is true for the various equations of state and phases of hadronic matter. Our analysis was for the case of a single equilibration channel, and it would be interesting to extend it to multiple channels, which may be relevant to both hadronic and quark matter (see appendix A of \\cite{Alford:2006gy} and Ref.~\\cite{Sa'd:2007ud}). In quark matter the non-Fermi liquid enhancement of the Urca rate \\cite{Schafer:2004jp} should further increase resonant effects. The application of our results to r-modes in neutron stars also raises interesting questions concerning the correct treatment of the crust \\cite{Steiner:2007rr}, and possible modification of the radial profile of the r-mode due to strong radial dependence of the bulk viscosity in layered stars such as hybrid stars.  Finally we note that at low temperatures the suprathermal enhancement of bulk viscosity becomes large, and the amplitude threshold for entering the suprathermal regime becomes low. This may not be relevant to the damping of r-modes because they are also damped by shear viscosity which becomes large at low temperature. But for other modes of compact stars, such as monopole pulsations, shear viscosity will not play such a significant role, and suprathermal bulk viscosity might be the dominant source of damping if external perturbations make the amplitude large enough. This might be relevant to old, cold, accreting stars in binary systems."
    },
    "1005/1005.0901_arXiv.txt": {
        "abstract": "The dependence of the long-term optical/UV variability on the spectral and the fundamental physical parameters for radio-quiet active galactic nuclei (AGNs) is investigated. The multi-epoch repeated photometric scanning data in the Stripe-82 region of the Sloan Digital Sky Survey (SDSS) are exploited for two comparative AGN samples (mostly quasars)  selected therein, a broad-line Seyfert\\,1 (BLS1) type sample and a narrow-line Seyfert\\,1 (NLS1) type AGN sample within redshifts 0.3--0.8. Their spectral parameters are derived from the SDSS spectroscopic data. It is found that on rest-frame timescales of several years the NLS1-type AGNs show systematically smaller variability compared to the BLS1-type. In fact, the variability amplitude is found to correlate, though only moderately, with the Eigenvector\\,1 parameters, i.e., the smaller the \\hb\\ linewidth, the weaker the [O\\,III] and the stronger the \\feii\\ emission, the smaller the variability amplitude is. Moreover, an interesting inverse correlation is found between the variability and the Eddington ratio, which is perhaps more fundamental. The previously known dependence of the variability on luminosity is not significant, and that on black hole mass---as claimed in recent papers and also present in our data---fades out when controlling for the  Eddington ratio in the correlation analysis, though these may be partly due to the limited ranges of luminosity and black hole mass of our samples. Our result strongly supports that an accretion disk is likely to play a major role in producing the opitcal/UV variability. ",
        "introduction": "Variation of optical/UV emission on timescales from hours to years is one of the defining characteristics of Active Galactic Nuclei (AGNs). Numerous previous studies of AGN/quasar ensembles have shown that the variability amplitude is dependent on wavelength  \\citep[e.g.,][]{cutr85, di96, helf01}, luminosity \\citep[e.g.,][]{hook94, give99, hawk00}, time lag \\citep[e.g.,][]{de05} and redshift \\citep[e.g.,][]{vand04}, and yet has a large scatter. While it is generally believed that the optical/UV light is emitted from an optically thick accretion disk powered by a central massive black hole, the physical processes that produce its  variability as observed are not understood yet. Interestingly, recent progresses link the variability with some fundamental parameters of AGN. A correlation of variability with \\mbh\\ was reported by \\citet{wold07}. Moreover, \\citet{wilh08} found correlations of the variability with both luminosity (inversely) and \\mbh, based on which the authors speculated that variability is inversely related to the Eddington ratio (\\redd).  Prominent line emission in the optical/UV band is another defining feature of AGNs. A set of strong correlations has been revealed to link the emission line properties, i.e.\\ a narrow \\hb\\ line associated with strong optical \\feii, weak \\oiii\\ \\citep[][]{boro92}. These correlations form the so-called  eigenvector 1 (E 1)  of AGN, which is suggestively driven by some underlying, more fundamental physical parameters, most likely the Eddington ratio \\citep[][]{sule00, boro02, dong09a}.  However, the  relationship between the above two AGN characteristics---the optical/UV variability and the emission line properties (or E1)---is poorly explored observationally so far. The same is true for the relation between the variability and the Eddington ratio. In this Letter, we report our investigation on this issue, as the first result of our comprehensive studies of AGN variability by making use of the valuable multi-epoch repeated photometric scanning of the Stripe-82 region in the Sloan Digital Sky Survey (SDSS). One virtue of our study improving upon previous ones is the inclusion of AGNs  with optical spectra characteristic of narrow-line Seyfert\\,1 (NLS1), in addition to typical broad-line AGNs (BLAGN) of the broad-line Seyfert\\,1 (BLS1) type, whose variability has been extensively studied. The NLS1-type AGNs extend to the extremes in the E1 space and to high \\redd\\ values. We use the $\\Lambda$-dominated cosmology with $H_{0} =$ 70\\,\\kmps Mpc$^{-1}$, $\\Omega_{m} = 0.3$, and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "It is intriguing that the optical/UV variability is closely related to the AGN E1 parameters. This can be understood in light of the same interesting correlations with the Eddington ratio, also found here, provided that \\redd\\ is indeed the underlying driver of the E1 correlations. Since the variability--\\redd\\ correlation is significant enough for the BLS1-type sample alone without including NLS1-type, our result does not rely on the \\mbh\\ (thus \\redd) estimation for NLS1-type AGNs, which is still somewhat controversial. The fact that the correlations of the variability with the E1 parameters and \\redd\\ are as strong as those among the E1 parameters and \\redd\\ themselves suggests that the E1 correlations may be extended to those involving the optical/UV variability.  The correlation between the variability and \\mbh, though also present in our data, turns out to be insignificant when \\redd\\ is taken into account. We thus suggest that such a correlation, as claimed in recent studies \\citep[][]{wold07, wilh08}, might be spurious and induced from the variability-\\redd\\ relation. The previously known variability--luminosity inverse correlation is found to be insignificant or at most weak in our data. This may be partly due to the fact that our analysis is capable of recovering AGNs (some are NLS1-type) with small variability in the low luminosity regime (see Figure\\,\\ref{correlation_mass_edd}) which were largely missed in previous work for various reasons. However, our morphology and redshift cuts restrict both the luminosity and \\mbh\\ ranges of our samples, which are important for comparisons with previous work. For example, \\citet{vand04} have shown that a factor of 100 change in luminosity corresponds to a factor of 4 change in variability. Thus it cannot be ruled out that the lack of significant dependence on either luminosity or \\mbh\\ may be due to their limited ranges in our samples.   The ensemble weaker variability of NLS1- compared to BLS1-type AGNs, as evidently found here and also suggested by \\citet{klim04} based on a small sample, is simply a natural consequence of the correlations with the E1 parameters and the Eddington ratio. The dependence on \\redd\\ can be understood qualitatively in terms of decreasing variability with increasing radius of the accretion disk, independent of the actual physical processes that drives the variability. Such radius-dependent variability is suggested by the fact that the variability amplitude increases as wavelength decreases \\citep[e.g.][]{vand04} in the context of thermal disk model. As the accretion rate $\\dot{m}$ (in units of the Eddington rate) increases, the continuum emission region that dominates a given waveband moves outward to a larger radius $r$ (in units of the Schwarzschild radius) because of $\\lambda\\propto T^{-1}\\propto (M/\\dot{m})^{1/4} r^{3/4}$; and hence decreasing variability is expected. We note that such a trend of optical variability seems at odds with that in the X-ray band \\citep{papa04}, though in the X-ray case more convincing results are needed. The different behavior may result from the different mechanisms and emitting regions of the radiation, as well as different variability mechanisms,  in the two bands \\citep[see e.g.][]{uttl06}.   Our finding of the strong correlations of the variability with the E1 parameters and \\redd\\ imposes interesting constraints to AGN variability models. Foremost, this indicates that the optical/UV variability must be intrinsic to AGN activity; any external scenarios, such as gravitational microlensing \\citep[][]{hawk00}, multiple supernovae and star collisions \\citep[][]{cid00}, can be ruled out as the dominant process. For intrinsic variability models, it is not clear whether the mechanism causing variability is some minor, secondary effect or related to the main energy generation process, as discussed in \\citet{gask08}. Our result favors the latter since the Eddington ratio is directly related to the mass accretion rate of AGN at a given accretion efficiency. The correlation with \\redd\\ as found here suggests that the accretion disk most likely plays a critical role in producing the observed opitcal/UV variability, since the disk structure is depending on the mass accretion rate. For instance, the model proposed by \\citet[][]{li08} that the variability is caused by change of accretion rates can reproduce {\\em qualitatively} the observed inverse variability--\\redd\\ relation (Li S., private communication)."
    },
    "1005/1005.5009_arXiv.txt": {
        "abstract": "MICADO is the adaptive optics imaging camera for the E-ELT. It has been designed and optimised to be mounted to the LGS-MCAO system MAORY, and will provide diffraction limited imaging over a wide ($\\sim$1\\,arcmin) field of view. For initial operations, it can also be used with its own simpler AO module that provides on-axis diffraction limited performance using natural guide stars. We discuss the instrument's key capabilities and expected performance, and show how the science drivers have shaped its design. We outline the technical concept, from the opto-mechanical design to operations and data processing. We describe the AO module, summarise the instrument performance, and indicate some possible future developments. ",
        "introduction": "MICADO is the {\\em Multi-AO Imaging Camera for Deep Observations}, designed to work with adaptive optics (AO) on the E-ELT. It has been optimised for the multi-conjugate adaptive optics (MCAO) module MAORY\\cite{dio10,fop10}; but it is also able to work with other adaptive optics systems, and includes a separate module to provide a single conjugate adaptive optics (SCAO) capability\\cite{cle10} using natural guide stars during early operations (see Section~\\ref{sec:ao}). As this simple AO mode sets low requirements on the telescope and facilities (e.g. no lasers are required), it is an optimum choice for demonstrating the scientific capabilities of the E-ELT at the earliest opportunity. The optical relay and support structure for SCAO provide the same opto-mechanical interface as MAORY, and in principle enable MICADO to be used with other AO systems such as ATLAS\\cite{fus10}. This phased approach means that MICADO will be able to make use of increasingly sophisticated AO systems as they become available.  MICADO is compact and is supported underneath the AO systems so that it rotates in a gravity invariant orientation. It is able to image, through a large number of selected wide and narrow-band near infrared filters, a large 53\\arcsec\\ field of view at the diffraction limit of the E-ELT. MICADO has two arms. The primary arm is a high throughput imaging camera with a single 3\\,mas pixel scale. This arm is designed with fixed mirrors for superior stability, thus optimizing astrometric precision. In addition, MICADO will have an auxiliary arm to provide an increased degree of flexibility. In the current design, this arm provides (i) a finer 1.5\\,mas pixel scale over a smaller field, and (ii) a 4\\,mas pixel scale for a simple, medium resolution, long-slit spectroscopic capability. In principle the auxiliary arm also opens the door to other options, including a `dual imager' based on a Fabry-Perot etalon to image emission line and continuum wavelengths simultaneously, coronography (perhaps implemented in a comparable way to that in NACO\\cite{boc04}), or a high time resolution detector.  \\begin{figure} \\begin{center} \\includegraphics[height=6cm]{micado-maory_zoom.eps} \\end{center} \\caption{Illustration of MICADO mounted under the MCAO system MAORY on the E-ELT Nasmyth platform. \\label{fig:photo} } \\end{figure} ",
        "conclusions": ""
    },
    "1005/1005.0289_arXiv.txt": {
        "abstract": "{}{}{}{}{} \\abstract {} {Our analysis is aimed at characterizing the properties of the integrated spectrum of active galactic nuclei (AGNs) such as the ubiquity of the Fe K$\\alpha$ emission in AGNs and the dependence of the spectral parameters on the X-ray luminosity and redshift.} {We selected 2646 point sources from the 2XMM catalogue at high galactic latitude ($|BII|$ $>$ 25 degrees) and with the sum of EPIC-PN and EPIC-MOS 0.2-12 keV counts greater than 1000. Redshifts were obtained for 916 sources from the Nasa's Extragalactic Database. We excluded sources classified as HII regions, groups/clusters, star-forming/starburst galaxies. The final sample consists of 507 AGN. Individual source spectra have been summed in the observed frame to compute the integrated spectra in different redshift and luminosity bins over the range 0$<$z$<$5. Detailed analysis of these spectra has been performed using appropriately normalized background spectra and exposure time weighted response and ancillary files.} {We find that the narrow Fe K$\\alpha$ line at 6.4 keV is significantly detected up to z=1. The line equivalent width decreases with increasing X-ray luminosity in the 2-10 keV band (``Iwasawa-Taniguchi effect''). The anticorrelation is characterized by the relation $\\log(EW_{Fe}) = (1.66\\pm0.09) + (-0.43\\pm0.07) \\log(L_{X,44})$, where $EW_{Fe}$ is the rest frame equivalent width of the neutral iron K$\\alpha$ line in eV and $L_{X,44}$ is the 2-10 keV X-ray luminosity in units of $10^{44}$ erg s$^{-1}$. The equivalent width is nearly independent of redshift up to z$\\sim$0.8 with an average value of 101$\\pm$40 (rms dispersion) eV  in the luminosity range 43.5$\\leq$logL$_X$$\\leq$44.5. Our analysis also confirmed the hardening of the spectral indices at low luminosities implying a dependence of obscuration on luminosity.} {We confirm that the neutral narrow Fe K$\\alpha$ line is an almost ubiquitous feature of AGNs. We find compelling evidence for the ``Iwasawa-Taniguchi effect'' over a redshift interval larger than probed in any previous study. We detect no evolution of the average rest frame \\eqw of \\ka line with redshift.} ",
        "introduction": "Active galactic nuclei (AGNs) are believed to be powered by accretion onto supermassive black holes (SMBH) \\citep{Salpeter1964, Lynden-Bell1969, Rees1984}. Most of the accretion power is released in the innermost region around the central SMBH \\citep{Shakura1973, Pringle1981}, where the emerging spectrum is significantly affected by strong relativistic effects \\citep{Page1974, Cunningham1975}. The Fe K$\\alpha$ emission line is the most prominent spectral feature observed in the X-ray spectra of AGNs. The line profile serves as a powerful diagnostic tool to probe its origin and provides unique information on the dynamics of the emitting region \\citep{Fabian1989, Laor1991}.  The Fe K$\\alpha$ emission was first identified as a common feature in the X-ray spectra of Seyfert galaxies observed by Ginga \\citep{Pounds1990, Nandra1994}. Early studies of type 1 AGNs with ASCA revealed resolved, broad and asymmetric Fe K$\\alpha$ lines in the majority of sources under study. The asymmetric lines were described as being originated from the accretion disc \\citep{Nandra1997}. Recent Chandra and XMM AGNs observations show that narrow Fe K emission is a ubiquitous feature in their X-ray spectra  \\citep{Yaqoob2004, Pounds2002}. However, unambiguous broad lines have been significantly detected and characterized only in a few Seyfert 1 galaxies, examples being MCG-6-30-15 \\citep{Fabian2002}, Mrk 205 \\citep{Reeves2001} and Mrk 509 \\citep{Pounds2001}. The number of spectral counts in AGNs is a major constraint in characterizing the line properties, since an accurate modeling of the underlying continuum is essential to measure the line parameters.  Several studies have been conducted to characterize the average Fe K$\\alpha$ emission in local and distant AGNs. \\citet{Nandra2007} performed a spectral analysis of a sample of 26 nearby Seyfert galaxies (z$<$0.05) observed by XMM-{\\it Newton}. They found evidence for broad line emission in around 65\\% of the sample with typical equivalent width (EW) being $\\sim$77$\\pm$16 eV, when fitted with a broad Gaussian. \\citet{Gu2006} performed a similar analysis using a sample of 102 AGNs (including Seyfert galaxies, Radio Quiet Quasars) observed with XMM-{\\it Newton} at z$<$0.5. They detected relativistic lines in about 25\\% of the sample objects ($\\simeq$50\\% when the ``well exposed'' spectra with $\\gtrsim$10000 net counts in the 2-10 keV band are considered). They also divided the sample in four equally populated luminosity classes and stacked residuals by modeling the continuum in the 2.5-15 keV energy range with an absorbed power law and Compton reflection component. They found that the average EWs derived from the stacked spectral residuals analysis are $\\lesssim$150 eV for all luminosity classes. They reported that the strongest relativistic profiles may be present in low-luminosity (L$_X$$<$10$^{43}$ erg s$^{-1}$) AGNs.  \\citet{Alina2005} derived an average rest frame spectrum of AGNs detected in a 770 ksec XMM-{\\it Newton} observation of the Lockman Hole field. Using a sample of 104 X-ray sources (53 type 1 and 41 type 2 AGNs and 10 galaxies) covering 0$<$z$<$4.5 in redshift space, they reported a clear relativistic line in the average rest frame spectra with an EW of $\\sim$560 eV and $\\sim$460 eV for the type 1 and type 2 AGNs, respectively. \\citet{Corral2008} computed the mean Fe emission from a large sample of more than 600 type 1 AGNs spanning a redshift range up to $\\sim$3.5. They detected significant unresolved Fe K$\\alpha$ emission line around 6.4 keV with an EW$\\sim$90 eV and found no compelling evidence of any significant broad relativistic emission line in the average spectrum.  \\citet{Brusa2005} studied the average spectra of AGNs detected in the Chandra Deep Field North and South. They used a large sample of 352 spectroscopically identified AGNs and stacked X-ray counts in the observed frame in seven redshift bins over the range 0.5$<$z$<$4. The measured Fe K$\\alpha$ line EWs are consistent with those reported by \\citet{Alina2005}. However, a broad component was not required by the data. Based on these recent studies, the ubiquity of relativistically broadened Fe K$\\alpha$ emission in AGNs and redshift dependence of the line strength are still unclear.  Studying the relation between emission line strength (characterized by the EW) and continuum luminosity is very important because it unveils the possible association between the emitting regions of these two spectral components. \\citet{Baldwin1977} first investigated the continuum luminosity dependence of the C IV $\\lambda$1549 EW and discovered that the EW of the C IV line decreased with increasing UV luminosity, a phenomenon now known as the ``Baldwin Effect''. Since that time, confirming its existence and exploring its origin and evolution has become the subject of several studies \\citep[e.g.][]{Dietrich2002, Green2001}. \\citet{IT1993} reported the anticorrelation between the Fe K emission line EW and X-ray luminosity, based on Ginga observations of AGNs. This relation was designated as ``X-ray Baldwin effect'' or ``Iwasawa Taniguchi effect'' (hereafter IT effect). This trend was confirmed by later ASCA observations of broad iron lines \\citep{Nandra1997}. In recent years, the existence of the IT effect for the narrow Fe K$\\alpha$ emission lines in {\\it type 1 AGNs} observed by XMM-{\\it Newton} and Chandra has been well established \\citep{Page2004, Zhou2005, Bianchi2007, Wu2009}.  The most fundamental problem, however, is that the physical causes for the ``UV/optical Baldwin effect'' and ``IT effect'' still remain unknown. Several models have been proposed to account for these phenomena. \\citet{Mush1984} suggested that the observed relation is due to a combination of decreasing ionization parameter and covering factor with increasing continuum luminosity. \\citet{Korista1998} pointed out that the Baldwin effect is driven by a softening of the ionizing continuum and increasing metallicity towards higher luminosities. Other studies propose that the underlying physical parameters driving the Baldwin effect may be the Eddington ratio \\citep{Baskin2004, Warner2004, Dong2009} and black hole mass \\citep{Xu2008}. \\citet{Nandra1997} attributed the IT effect to the presence of an ionized skin on the accretion disc, with the degree of ionization increasing with luminosity \\citep{NayaB2000, NayaA2000}.  In the present study, we exploit the 2XMM catalogue to characterize the properties of the integrated spectrum to address the following key points : (i) ubiquity of the Fe K$\\alpha$ emission line in AGNs, (ii) origin of the IT effect and (iii) evolution of the Fe K$\\alpha$ \\eqw with redshift. The conventional approach to improve the counting statistics in AGNs involves unfolding the instrumental response by appropriately modeling the continuum and then shifting and adding the individual spectra, which in turn, maximizes the S/N ratio around the iron line energy. However, the main drawback is a model dependent parameterization of the underlying continuum. We adopt an alternative approach, first implemented by \\citet{Brusa2005}, in which source spectra are stacked in narrow redshift intervals.  The paper is organized as follows. In the following Section, the selection criteria and properties of the sample are discussed. Section 3 outlines the spectral analysis. The following Sections then present and discuss the results. Thoughout this work a cosmology with $\\Omega_m$=0.27, $\\Omega_\\Lambda$=0.73 and H$_0$=70 km s$^{-1}$ Mpc$^{-1}$ is used. ",
        "conclusions": "We have compiled a large sample of 507 sources selected from the 2XMM catalogue covering the redshift range 0$<$z$<$5 and nearly 6 orders of magnitude in the X-ray luminosity. To improve the signal-to-noise ratio and investigate the average X-ray spectral properties of the sources, we computed the integrated spectra in narrow intervals of redshift and luminosity. We performed simulations to assess the accuracy of the stacking procedure. The main results can be summarized as follows. \\begin{itemize} \\item[$\\bullet$] {We confirm that the narrow Fe K$\\alpha$ line is an almost ubiquitous feature in the X-ray spectra of AGNs.} \\item[$\\bullet$] {In the integrated spectra spanning redshift range 0$<$z$<$1.3 at nearly constant luminosity we detect no strong trend in the Fe K$\\alpha$ line EW with redshift.} \\item[$\\bullet$]{We find a significant anticorrelation between the Fe K$\\alpha$ EW and 2-10 keV rest frame luminosity across a wide redshift range not probed in prior studies. The slope of the anticorrelation we obtained is steeper than the value reported in previous studies, likely due to the difference in composition of the samples. The exact physical cause of the IT effect is unknown, one promising explanation is a decline in the covering factor of the putative molecular torus as the luminosity increases.} \\item[$\\bullet$]{We find a hardening of the spectral indices at low luminosities in the redshift bin 0.0-0.2. This can be linked to the dependence of obscuration with luminosity.} \\end{itemize}  Large area X-ray surveys from sensitive next generation satellites (e.g., eROSITA and WFXT), when coupled with deep multiwavelength photometry (e.g., PanSTARRS and LSST) and spectroscopic campaigns (SDSS, LAMOST, and Euclid) will allow to study the dependence of the EW strength with redshift and luminosities on samples of several order of magnitudes larger number statistics providing a unique sampling of the L-z plane.  \\acknowledgement Our results are based on observations obtained with XMM-{\\it Newton}, an ESA science mission with instruments and contributions directly funded by ESA member states and the USA (NASA). The XMM-{\\it Newton} project is supported by the Bundesministerium f\\\"ur Wirtschaft und Technologie/Deutsches Zentrum f\\\"ur Luft und Raumfahrt (BMWI/DLR, FKZ 50 OX 0001), the Max-Planck Society and the Heidenhain-Stiftung and also by PPARC, CEA, CNES, and ASI. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. GH acknowledges support from the German Deutsche Forschungsgemeinschaft, DFG, Leibniz Prize (FKZ HA 1850/28\u20131). This work is partially supported by ASI-INAF grants I/023/05/00 and I/088/06. This research was also supported by the DFG cluster of excellence ``Origin and Structure of the Universe\". The authors would like to thank the anonymous referee for the encouraging report and useful comments, which improved the paper significantly. We also thank Stefano Bianchi and Kazushi Iwasawa for useful discussions."
    },
    "1005/1005.5653_arXiv.txt": {
        "abstract": "{We report the detection of absorption lines by the reactive ions \\OHp, \\HHOp and \\HHHOp\\ along the line of sight to the submillimeter continuum source G10.6$-$0.4 (W31C). We  used the \\Herschel\\ HIFI instrument in dual beam switch mode to observe the ground state rotational transitions of \\OHp\\ at 971~GHz,  \\HHOp\\ at 1115 and 607~GHz, and \\HHHOp\\ at 984 GHz. The resultant spectra show deep absorption over a broad velocity range that originates in the interstellar matter along the line of sight to  G10.6$-$0.4 as well as in the molecular gas directly associated with that source. The \\OHp\\ spectrum reaches saturation over most velocities corresponding to the foreground gas, while the opacity of the \\HHOp\\ lines remains lower than 1 in the same velocity range, and the \\HHHOp line shows only weak absorption. For Local Standard of Rest (LSR) velocities between 7 and 50 kms$^{-1}$ we estimate total column densities  of $N$(\\OHp) $\\geq 2.5 \\times 10^{14}$ cm$^{-2}$, $N$(\\HHOp) $\\sim 6 \\times 10^{13}$ cm$^{-2}$ and $N$(\\HHHOp) $\\sim 4.0 \\times 10^{13}$  cm$^{-2}$. These detections confirm the role of O$^+$ and OH$^+$ in initiating the oxygen chemistry in diffuse molecular gas \\referee{and strengthen our understanding of the gas phase production of water}. The high ratio of the \\OHp\\ by the \\HHOp\\ column density implies that these species predominantly  trace low-density gas with a small fraction of hydrogen in molecular form.  } ",
        "introduction": "Oxygen chemistry in the interstellar medium is initiated by cosmic-ray or X-ray ionization of hydrogen. The resulting H$^+$ and H$_3^+$ transfer their charge to form the O$^+$ and OH$^+$ ions, both of which react rapidly with H$_2$ molecules. Subsequent H-atom abstraction reactions proceed to the terminal ion H$_3$O$^+$ with high efficiency, unless the electron fraction is unusually high. H$_3$O$^+$ does not react with H$_2$ or any other abundant neutral species, but is destroyed by dissociative recombination with electrons. Recombination produces O, OH, and H$_2$O. \\referee{This sequence is the main gas-phase chemical route to water in the interstellar medium.} Until now, this classical scheme of the oxygen chemistry (Herbst \\& Klemperer 1973) has been tested mainly through observations of OH and H$_2$O, whose abundances are affected by many other processes as well, and by a few observations of H$_3$O$^+$. With the HIFI instrument on the \\Herschel \\ satellite it is now possible to observe the strongest transitions of light hydrides, including the reactive molecular ions OH$^+$ and H$_2$O$^+$ that are thought to play crucial intermediate roles in the formation of OH and H$_2$O.  \\referee{Previous attempts to observe interstellar  H$_2$O$^+$ and OH$^+$ from optical spectroscopy are very limited. Smith et al. (1984) report an upper limit of $10^{12}$ cm$^{-2}$ for  H$_2$O$^+$ based on optical absorption spectroscopy. No observation of the OH$^+$ transition at  358.4 nm has been reported until very recently (Krelowski et al., private communication). Submillimeter telescopes have been more successful in detecting these ions.} With the Atacama Pathfinder EXperiment (APEX) telescope Wyrowski et al. (2010) detected two hyperfine-structure components of the OH$^+$ $N=1-0$ $J=0-1$ transition near 909.2 GHz towards the strong continuum source Sgr B2(M) and concluded that \\OHp \\ plays a major role in the chemistry of the diffuse molecular gas. Using HIFI on \\Herschel,  Ossenkopf et al. (2010) report the detection of o-\\HHOp \\ absorption in the three massive star-forming regions NGC~6334, Sgr~B2, and DR~21.  We present observations toward the massive star-forming region G10.6$-$0.4 (W~31C) located at $RA=18h10m28.70s$, $Dec =-19^\\circ55'50 s$ (J2000). The submillimeter-wave continuum emission is associated with a cluster of ultra-compact H\\,{\\sc ii} regions located in the so-called ``30 km s$^{-1}$ arm'' at a kinematic distance of 4.8 kpc (Fish et al. 2003, Corbel \\&\\ Eikenberry 2004). The molecular cloud surrounding the star-forming complex shows evidence of continuing accretion onto the embedded stars (Keto \\&\\ Wood 2006; Sollins \\&\\ Ho 2005). As discussed by Neufeld et al. (2010), line absorption by foreground gas has been detected from [OI] (Keene et al. 1999), $^{13}$CH$^+$ (Falgarone et al. 2005), HCO$^+$, HCN, HNC, CN (Godard et al. 2010), and atomic hydrogen (Fish et al. 2003). While the gas associated directly with G10.6$-$0.4 is detected in the velocity range $-5$ to $+5$ km s$^{-1}$, the foreground gas along the line of sight is confined to the velocity interval from $\\sim 5$ to 45 km s$^{-1}$. ",
        "conclusions": "\\referee{The detection of interstellar \\OHp, \\HHOp, and \\HHHOp provides a strong support to our understanding of the gas phase oxygen chemistry, in particular the gas phase route to water vapor. The combined analysis of the \\OHp\\ and \\HHOp\\  ground state} spectra has shown that they mostly reside in warmish atomic gas where they are  produced in a reaction sequence initiated by atomic H ionization followed by rapid reaction with traces of H$_2$. Their abundances and abundance ratio can be used to constrain the fraction of hydrogen in molecular form, and to set a lower limit for the cosmic ray ionization in the diffuse gas, where few other diagnostics are applicable. As these reactive molecular ions can be observed from space with the \\Herschel \\ Space Observatory, and in exceptional conditions from the ground, these preliminary conclusions should be verified by more extensive observations. It is also expected that strong features from both \\OHp \\ and \\HHOp\\ will be detected in the spectra of distant galaxies, where elevated fluxes of ionizing radiation are produced by active nuclei and intense star-formation activity."
    },
    "1005/1005.1714.txt": {
        "abstract": "We present a study on the impact of molecular outflows in the Perseus molecular cloud complex using the COMPLETE survey large-scale $^{12}$CO(1-0) and $^{13}$CO(1-0) maps. We used three-dimensional isosurface models generated in RA-DEC-Velocity space to visualize the maps. This rendering of the molecular line data allowed for a rapid and efficient way to search for molecular outflows over a large ($\\sim 16$ deg$^2$) area. Our outflow-searching technique detected previously known molecular outflows as well as new candidate outflows. Most of these new outflow-related high-velocity features lie in regions that have been poorly studied before. These new outflow candidates more than double the amount of outflow mass, momentum, and kinetic energy  in the Perseus cloud complex. Our results indicate that outflows have significant impact on the environment immediately surrounding localized regions of active star formation, but lack the energy needed to feed the observed turbulence in the {\\it entire} Perseus complex. This implies that other energy sources, in addition to protostellar outflows, are responsible for turbulence on a global cloud scale in Perseus. We studied the impact of outflows in six regions with active star formation within Perseus of sizes in the range of 1 to 4 pc. We find that outflows have enough power to maintain the turbulence in these regions and enough momentum to disperse and unbind some mass from them. We found no correlation between outflow strength and star formation efficiency for the six different regions we studied, contrary to results of recent numerical simulations. The low fraction of gas that potentially could be ejected due to outflows suggests that additional mechanisms other than cloud dispersal by outflows are needed to explain low star formation efficiencies in clusters.   %We present a study on the impact of molecular outflows in the Perseus molecular cloud complex using the COMPLETE survey large-scale $^{12}$CO(1-0) and $^{13}$CO(1-0) maps. We used   three-dimensional isosurface models generated in RA-DEC-Velocity space to visualize the maps. This  rendering of  the molecular line data allowed for a rapid and efficient way to search for molecular outflows over a large ($\\sim 16$ deg$^2$) area. Our outflow-searching technique detected previously known molecular outflows as well as new candidate outflows. Most of these new outflow-related high-velocity features lie regions that have been poorly studied before(i.e., in the outskirts of, and areas between, active star-forming regions). These new outflow candidates approximately double the amount of outflow mass, momentum, and kinetic energy associated with outflows in the Perseus molecular cloud complex, which show  the  importance of observing the whole molecular cloud complex in order to obtain a complete census of the outflows in the cloud. Our results indicate that outflows have considerable impact on the environment immediately surrounding  localized regions of active star formation, but lack the energy needed to feed the observed turbulence in the entire Perseus cloud complex, even with the increase in outflowing material found by our survey. This implies that other  energy sources, in addition to collimated outflows from protostars,  are responsible for turbulence on a global cloud scale in the Perseus complex.  We studied the impact of outflows in six regions with active star formation within the cloud complex of sizes in the range of 1 to 4 pc. We find that outflows have enough power to maintain, or substantially contribute to, the turbulence in their host regions and have enough energy and momentum to disperse and unbind some mass from them. We found no correlation between outflow strength and star formation efficiency for the six different clumps we studied, contrary to results of recent numerical simulations. In most star-forming clumps the fraction of gas that potentially could be ejected due to  outflows is   lower than that predicted by theoretical models and implied by observations of the star formation efficiency in clusters. We suggest that other mechanisms of cloud dispersal,  along with collimated outflows from young stars,  are needed in protostellar clusters. ",
        "introduction": "\\label{introsection}  Outflows are an intrinsic part of the star formation process, as all stars, both low and high mass, go through a mass-loss phase during their protostellar stages (e.g., Arce et al.~2007). The outflowing supersonic wind from a protostar can accelerate the surrounding molecular gas to velocities significantly greater than those of the quiescent cloud gas thereby producing a molecular outflow. Early on in the study of outflows from young stars it was realized that they have the potential to have a major effect on the dynamics and structure of their parent cloud \\citep{norman80}.  More recently, it has been suggested that in regions of low-mass star formation, outflows could be the leading disruptive agent that limits the lifetime of their parent molecular cloud \\citep{hartmann01}. Analytical and numerical studies indicate that outflows can couple strongly to the cloud and are highly efficient at driving turbulent motions (Matzner 2007; Nakamura \\& Li 2007; Cunningham et al. 2009; Carroll et al.~2009) and can also regulate the cloud's star formation efficiency (e.g.,  Matzner \\& McKee 2000;  Nakamura \\& Li 2007). Yet, other numerical studies suggest that protostellar outflows do a poor job at driving cloud turbulence, but can disrupt dense clumps and affect the cloud structure (Banerjee et al. 2007). These studies show the increased attention that research on the impact of outflows has obtained, as well as the need for targeted observations required to constrain the models and to reconcile their discrepancies.  The underlying physics in studies of outflow-environment interactions is mostly understood (even if difficult to model), however observations are crucial for constraining the various assumptions made in these models (e.g., outflow power, degree of collimation, mass outflow rate, outflow lifetime, etc.).  Significant constraints on the values of outflow characteristics require systematic observations of large outflow samples. Different physical processes triggered by protostellar outflows may be traced at different wavelengths (see, e.g., Hartigan et al.~2000). Shock-excited infrared lines of H$_2$ and optical Herbig-Haro (HH) objects trace the recently shocked gas that cools within a few years. On the other hand, the high velocity CO  trace the surrounding molecular gas entrained by the protostellar wind and remains visible by collisional excitation for much more time than the shocked gas. This is why CO has mostly been used for studying the impact of outflows on the surrounding molecular cloud. However, the CO  fails as an outflow tracer  in regions where there is little molecular gas (i.e., in the outskirts of the cloud), and so molecular outflows give a lower limit on the outflow momentum and energy injection into their (low-density) surroundings.   In these low-density regions other tracers like the HI (21~cm) and  CII (157 $\\micron$) lines should be useful for tracing the outflow.  %In the research presented here, we aim to shed light on this issue by conducting %a molecular outflow survey of the Perseus molecular cloud complex and conducting a %study of the impact of outflows over the entire molecular cloud complex.  Observational studies have shown that outflows, even from low-mass stars, can have an impact on their cloud at different distances from the source ranging from a few thousand AU to several parsecs.  Survey studies of the circumstellar gas within $10^4$ AU of low-mass protostars indicate outflows contribute significantly to the mass-loss of the surrounding dense gas \\citep{fuller02,arce06}.  An outflow's impact on its parent core, at distances of about 0.1 to 0.3~pc from the forming star, is evidenced through the detection of outflow-blown cavities and fast-moving (outflowing) dense gas  \\citep[e.g.,][]{tafalla07}.  Giant outflows from young stars with sizes exceding 1~pc in length are common \\citep{reipurth97,stanke00}.  These outflows can interact with the surrounding medium and induce changes in the velocity and density distribution of the parent cloud's gas at large distances from the source \\citep{arce01b,arce02a}.  Millimeter studies show that many molecular outflows produced by a group or cluster of young stars interact with a substantial volume of the cluster's environment, may sweep up the gas and dust into ``shells'' and can have the energy required for driving the turbulence in the cluster gas \\citep[e.g.,][]{knee00,williams03,stanke07}.  All these observational studies suggest that individual, groups, and clusters of outflows have a significant impact on their surroundings within a few parsecs from their sources.  The collective impact of all the outflows on an entire molecular cloud or on a molecular cloud complex (with size $\\sim $10~pc) in which they reside is, however, still unclear.  %\\vspace{-4in} \\begin{figure*} \\epsscale{1.3} %\\plotone{f1.eps} \\plotone{f1.pdf} \\vspace{-3in} \\caption{$^{13}$CO(1-0) integrated intensity map of the  Perseus molecular cloud complex from the COMPLETE survey (Ridge et al.~2006a). The six areas visualized in 3D Slicer are labeled I through VI. The black (solid line) boxes show the approximate boundaries of the star-forming regions used in our analysis in section~\\ref{analysis} (see also Tables~\\ref{regiontab} and \\ref{physregtab}). Note that the $^{13}$CO map is not corrected for the FCRAO beam efficiency. \\label{perseus_map}} \\end{figure*}  Large-scale (unbiased) molecular outflow surveys in the 1980's were crucial in estimating the molecular outflow energy in whole clouds, and acknowledging that outflows can have a significant impact on their host cloud.  These surveys were conducted toward clouds with high-mass star formation ---i.e., Mon OB1 \\citep{margulis86} and the Orion southern cloud \\citep{fukui86}--- using small (4 to 5~m) millimeter telescopes resulting in large beams (2.3\\arcmin \\/ to 2.7\\arcmin).  Only regions with detectable high-velocity gas ($v > 10$~km~s$^{-1}$) were then mapped at higher (sub-arcminute) resolution.  Beam dilution should have hampered the ability to detect small ( $< 0.4$~pc) and slow ($v < 10$~km~s$^{-1}$) outflows in these surveys and many such outflows were probably missed.  Comparing recent higher angular resolution observations of small regions included within the large-scale surveys conducted about two decades ago, it is clear that these surveys detected only the most powerful outflows and that individual outflows were not resolved in high-density regions; compare, for example, the results of \\citet{yu00} and \\citet{fukui86}. %(compare, for example, the results from  \\citet{fukui86} and \\citet{yu00}.  Most of the recent outflow studies are restricted to small fields, focusing on individual objects or small dense regions of known star formation activity, thereby providing a limited view of the outflows' impact on the environment at large (cloud) scales.  These biases inhibit a full census of the outflows in a cloud, prevent the identification of multi-parsec outflows, and limit the study of the outflows' impact on the entire  molecular cloud complex.  For example, the studies of \\citet{hat07} and \\citet{hat09} present a search for outflows towards known cores in the Perseus molecular cloud complex, and the study succeeds in deriving outflow properties of outflows in the Perseus cloud in a consistent way and finding a few new (previously undetected) outflows.  However, they fail to provide an unbiased survey of outflows across the entire Perseus cloud complex as most observations are concentrated within  $\\sim0.15$~pc from known cores and regions of active star formation. In contrast, the outflow survey presented here covers the full extent of the Perseus molecular cloud complex, and has the advantage of being able to provide information on how individual and groups of outflows effect the dynamics of the gas in the entire cloud complex and how they interact with their surroundings at different distances from the driving source (from about 0.06~pc to a few parsecs).  %\\subsection{The Perseus Molecular Cloud}  The Perseus molecular cloud complex is a chain of clouds with a total mass of a few thousand solar masses, and encompasses a total area of about 70~pc$^{2}$ \\citep{evans09}. There is a span of distance estimates to Perseus that range from 230~pc \\citep{cer90} to 350~pc \\citep{her83}, most probably because different regions of the cloud complex are at different distances along the line of sight, and thus using a single distance for the entire cloud may be inappropriate.  However, accurate distance estimates require maser parallaxes with high accuracy (e.g., Hirota et al.~2008), which do not yet exist for the entire cloud. To simplify our calculations, we adopt a single fiducial distance of  250$\\pm$50~pc, similar to other recent cloud-wide studies of Persues  \\citep{eno06, jor06, rid06data, reb07, evans09}, and caution that there may be significant differences in the distance  to different parts of the cloud. Perseus contains two rich protostellar clusters,  IC~348 and NGC~1333, and a number of other regions of active star formation including  B5, B1, L1448, and L1455 (see Figure~\\ref{perseus_map}).  Surveys in the infrared, sub-millimeter, and millimeter reveal a large population of low-mass pre-main sequence stars, embedded protostars, and starless cores in Perseus \\citep{lad93,aspin94,lad95, hat05, eno06, kirk06, muench07, guter08, evans09}. Only one B5 star in IC~348 (HD~281159) is confirmed to reside in the Perseus cloud, but there might be a few other high-mass stars that interact with the cloud (through their winds and/or UV radiation) even though they were not necessarily formed in the cloud complex (see, e.g., Walawender et al.~2004; Ridge et al.~2006b; Kirk et al.~2006; Rebull et al.~2007). There is also a large number of nebulous objects associated with outflow shocks (i.e., HH objects and H$_2$ knots ) that have been identified in the cloud complex \\citep{bal96a,bal97,yan98,wal05, davis08}.  The whole Perseus region was first surveyed in $^{12}$CO by \\citet{sar79}, and since then has been mapped in CO at different angular resolutions (all with beams $>1\\arcmin$) by a number of other authors \\citep[e.g.,][]{bachiller86,ungerechts87,padoan99,sun06}. These maps show a clear velocity gradient in the Perseus molecular cloud complex where the central cloud (LSR) velocity increases from about 4.5 km~s$^{-1}$ at the western edge of the cloud to about 10~km~s$^{-1}$ at the eastern end. The large velocity gradient in the gas across the entire complex and the fact that different parts of the Perseus cloud appear to have different distances (see above) could possibly indicate that the complex is made up of a superposition of different entities. Recently, the Perseus molecular cloud complex was also observed (and studied) in its entirety in the mid- and far-infrared as part of the ``From Molecular Cores to Planet-forming Disks'' (ak.a., c2d) {\\it Spitzer} Legacy Project \\citep{jor06,reb07, evans09}.          % Move pieces to other parts of the paper if necessary: {\\it Visualizing the velocity structure of a cloud in 3D has the advantages of being easier to understand and interpret as compared to channel maps.  Traditionally, one has to look through the channel maps one by one (usually in a movie format) or use position-velocity diagrams to construct a 3D interpretation.  The creation of a 3D model in essence displays all the maps at the same time by stacking them all together and drawing a 3D surface around the emission.  As will be discussed in Section 3, physical phenomena like outflows are easily detectable in the form of spikes.  Conventionally, outflows are detected by either looking through channel maps for a blue shifted outflow lobe at a low velocity range as compared to the local standard of rest velocity (LSR), and a red shifted lobe at a higher velocity range as compared to the LSR.}  3D modeling lets you see all the velocity values at the same time so blue and red shifted lobes are concurrently visible.  {\\it The other way to find outflows is to individually examine known young stars in a region and observe them to see if they have an outflow.}  This is inefficient since not all young stars have outflows.  The 3D modeling used in this paper proved to be an excellent and efficient way to identify outflows, particularly over a large area such as Perseus which is 8.7 square degrees.  \\ ",
        "conclusions": "A survey of all the high velocity outflow features was conducted for Perseus using $^{12}$CO and $^{13}$CO molecular line maps from the COMPLETE Survey in order to gain a better understanding of how outflows effect their host cloud complex over large areas. Our $\\sim 6.25\\arcdeg \\times 3\\arcdeg$  maps cover almost the entire Perseus molecular cloud complex which allows us to search for outflows throughout the extent of the cloud and not just in the immediate surrounding of active star-forming regions as in most previous studies. We utilize a novel way to search for outflows using 3D visualization of the molecular cloud in position-position-velocity space. Rendering the molecular line data this way  shows high-velocity blobs as clearly visible spikes that stick out from the general cloud gas distribution. Using this technique we detected previously known outflows and new outflow candidates. Our survey results show that most molecular outflows in Perseus concentrate close to a few groups and clusters of protostars,  and there are vast areas of the cloud where there is little or no molecular outflow activity. We compiled a list of 60 COMPLETE Perseus Outflow Candidates (CPOCs) that lie across the entirety of Perseus, with most of them found around the periphery of, and in between, active star forming regions. We identified a few of these CPOCs as being previously unknown extensions of  previously known outflows. In particular, we find new extensions to the well-known  B5-IRS1 outflow which now has a total projected length of $\\sim 1\\arcdeg$ (or about 4.4 pc).  The newly identified candidate outflows more than double the amount of molecular outflow mass, momentum and kinetic energy in Perseus. Outflows have considerable impact on the environment immediately surrounding localized regions of active star formation, but lack the energy needed to feed the observed turbulence in the entire Perseus cloud complex, even with the increase in outflowing material found by our study. We propose there must be another energy source, in addition to collimated outflows from protostar, responsible for maintaining the turbulence on a global cloud scale. We leave the discussion on spherical winds from young stars as a possible source of turbulence in Perseus for a subsequent paper.  We compared the total outflow momentum, energy  and power in the different local star-forming regions ---areas of active star formation with numerous outflows and sizes 1 to 4 pc---  with the energetics of the host cloud sub-region to quantitatively assess the outflow impact on their surroundings. Our results indicate that outflows have enough power (or can contribute substantially) to maintain the turbulence in their local environment.  The numerical simulations  of Nakamura \\& Li (2007) show that outflow-driven turbulence can affect the star formation efficiency in their parent cluster-forming clump and we use our data (together with data from the literature) to investigate if this is the case in Perseus. We detect no correlation between star formation efficiency and outflow strength, but the large uncertainties in these values prevent us from reaching any definitive conclusion on this issue.  Our quantitative assessment of the impact of current  outflows on Perseus indicates that outflows have the potential to disperse and unbind some mass from the cloud. The current outflow energy and the amount of material that current outflows can potentially unbind, compared to the total mass of the region, is not enough to cause major disruptions to the cloud. The total outflow momentum injection throughout the entire lifetime of the cloud would have to be more than an order of magnitude greater than the detected current outflow momentum in order for gas dispersion by outflows to cause a serious impact to the integrity of their host cloud. We argue that, in addition to collimated outflows from young stars, other mechanisms of cloud dispersal are needed to explain the low star formation efficiency in clusters.  %predicted %to produce the predicted (30\\% to 50\\%, Matzner \\& McKee~2000) % and observed  %Only if  the %outflow and star formation activity in the cloud is constant (or increases with time) throughout the lifetime of the cloud, and % outflows inject more than an order of magnitude more momentum into their surroundings than our estimates of the % current observed outflow momentum  %then it may be that outflows could inject an order of magnitude more momentum into their surroundings than what we estimate using the current observed outflows.  %Even if this extrapolation is correct, in most cases the fraction %of gas that could be ejected from cluster-forming clumps due to  outflows is %lower than that predicted by theoretical models and implied by observations of the star formation efficiency in clusters."
    },
    "1005/1005.0847_arXiv.txt": {
        "abstract": "We extend our maximum likelihood method for reconstructing the cluster-mass cross-correlation from cosmic microwave background (CMB) temperature anisotropies and develop new estimators that utilize six different quadratic combinations of CMB temperature and polarization fields. Our maximum likelihood estimators are constructed with delensed CMB temperature and polarization fields by using an assumed model of the convergence field and they can be iteratively applied to a set of clusters, approaching to the optimal condition for the lensing reconstruction as the assumed initial model is refined. Using smoothed particle hydrodynamics simulations, we create a catalog of realistic clusters obtainable from the current Sunyaev-Zel'dovich (SZ) surveys, and we demonstrate the ability of the maximum likelihood estimators to reconstruct the cluster-mass cross-correlation from the massive clusters. The iTT temperature estimator provides a signal-to-noise ratio of a factor 3 larger than the iEB polarization estimator, unless the detector noise for measuring polarization anisotropies is controlled under $3\\muK$. ",
        "introduction": "In the past few years we have seen rapid development in the measurements of cosmic microwave background (CMB) anisotropies on arcminute scales. Higher precision measurements of CMB temperature anisotropies are available from the completed missions such as the Arcminute Cosmology Bolometer Array Receiver \\cite{READET09} and the Cosmic Background Imager \\cite{SIMAET09} and the ongoing experiments with better angular resolution and multi-frequency channels such as the Atacama Cosmology Telescope (ACT, \\cite{HIACET09}) and the South Pole Telescope (SPT, \\cite{STADET08}). While most of the current polarization experiments are aimed at measuring CMB polarization anisotropies on large scales, ACT and SPT will be capable of measuring CMB polarization anisotropies on arcminute scales.  Especially, ACT and SPT are designed for measuring the Sunyaev-Zel'dovich (SZ) effect arising from the Thomson scattering off hot electrons in clusters and thereby detecting massive clusters. These SZ surveys can explore the growth of structure by measuring the abundance of massive clusters and the expansion history of the universe by measuring their correlation functions to probe the evolution of dark energy (see, e.g., \\cite{DETF06}). With precisely measured distance and their omnipresence behind all the clusters, CMB anisotropies can be used as a distant background source for weak lensing measurements, providing the cluster-mass cross-correlation and complementing the measurements of cluster mass and their abundance from the same SZ surveys. Moreover, weak lensing measurements of the CMB can be combined with galaxy weak lensing measurements of the same clusters from optical follow-up surveys to measure the source distance ratios in a model-independent way, constraining the expansion history of the universe \\cite{HUHOVA07}.  Weak lensing of the CMB by large-scale structure is efficiently probed by using quadratic estimators, measuring the deviation of the correlation function of CMB anisotropies from the otherwise statistically isotropic correlation \\cite{HU01b}. However, it was shown that the standard method of using quadratic estimators is compromised \\cite{MABAMEET05} on cluster scales and needs additional free parameters \\cite{HUDEVA07} for calibration, as this method rests on the linear approximation in the lensing effect which breaks down near massive clusters. This problem was tackled by delensing the observed CMB temperature field and analyzing the likelihood of the delensed fields \\cite{YOZA08}.  We showed that if the assumed initial model for delensing is a good approximation to the true underlying matter distribution, our new estimator based on the delensed CMB temperature fields becomes an optimal estimator and it can be iteratively applied to given measurements of CMB temperature anisotropies, until the assumed initial model converges to the true matter distribution and the likelihood is maximized. Here we extend our maximum likelihood estimation method for reconstructing the cluster-mass cross-correlation with CMB temperature anisotropies and apply it to CMB polarization anisotropies, yielding new maximum likelihood estimators with six different quadratic combinations of CMB temperature and polarization fields. We demonstrate their applicability by using realistic clusters that can be found in the current CMB experiments.  The rest of the paper is organized as follows. We briefly review the formalism for weak lensing of CMB anisotropies and derive our maximum likelihood estimators in Sec.~\\ref{sec:for}. In Sec.~\\ref{sec:sim} we describe our numerical simulations to model realistic clusters and we construct a catalog of massive clusters with CMB secondary anisotropies. We reconstruct the cluster-mass cross-correlation by using our maximum likelihood estimators and quantify their signal-to-noise ratios in Sec.~\\ref{ssec:iqe}, and we investigate the impact of the kSZ contamination in Sec.~\\ref{ssec:ksz}. We summarize our findings and conclude in Sec.~\\ref{sec:dis}. For illustrative purposes, we adopt a flat $\\Lambda$CDM universe with cosmological parameters ($\\OM h^2=0.127$, $\\OB h^2=0.0222$, $h=0.73$, $n_s=0.95$, $\\rms=0.78$), consistent with the recent estimation (e.g., \\citep{TESTET06,KODUET08}). ",
        "conclusions": "\\label{sec:dis} We have generalized the maximum likelihood estimators \\cite{YOZA08} for reconstructing the cluster-mass cross-correlation to utilize delensed CMB polarization fields, and we have tested our new maximum likelihood estimators for the lensing reconstruction against numerical simulations. The observed CMB temperature and polarization fields are delensed based on an initial model of the convergence field and our maximum likelihood estimators (called improved quadratic estimators) provide improvements over the assumed initial model. This process can be iterated until numerical convergence is achieved. Compared to the standard and modified quadratic estimators \\cite{MABAMEET05,HUDEVA07}, our improved quadratic estimators have no free parameters and provide unbiased reconstructions even in the regime where the linear approximation in the lensing effect breaks down.  We have adopted smoothed hydrodynamic simulations to model realistic clusters, and the improved quadratic estimators that use six different combinations of CMB temperature and polarization fields can reconstruct the underlying matter distribution in a non-parametric way to obtain the cluster-mass cross-correlation. Their ability to reconstruct the cluster-mass cross-correlation is determined by the ratio of the intrinsic CMB power spectrum to the detector noise power spectrum on cluster scales. Given the same experimental specifications for measuring CMB temperature and polarization anisotropies, iTT estimators can reconstruct the cluster-mass cross-correlation by a factor of~3 in the signal-to-noise ratio, better than iEB estimators, which are better suited for reconstructing large-scale structure.  For the gas mass fraction $f_\\up{gas}=0.117$, the numerical simulations show that the kinetic Sunyaev-Zel'dovich (kSZ) contamination is $\\Delta\\Tksz\\simeq30\\muK$ at the innermost region of clusters. As clusters are more compact at high redshift with weak redshift dependence of their peculiar velocity, the spectrally indistinguishable kSZ contamination in CMB temperature anisotropies poses a significant challenge to the lensing reconstruction using iTT estimators. However, polarization estimators such as iEB estimators or more practically iEE estimators can be used for the cross-check of the lensing reconstruction. Especially, these polarization estimators are more desirable given that there is relatively little kSZ contamination in CMB polarization anisotropies and they perform better than iTT estimators in an ideal CMB experiment with no detector noise and telescope beam. However, with the prospect of CMB polarization experiments competitive with the current CMB temperature experiments far in the future, masking the central region of clusters to mitigate the kSZ contamination will be needed to take full advantage of precise CMB temperature measurements in the lensing reconstruction."
    },
    "1005/1005.3655_arXiv.txt": {
        "abstract": "While Bayesian model selection is a useful tool to discriminate between competing cosmological models, it only gives a relative rather than an absolute measure of how good a model is. Bayesian doubt introduces an unknown benchmark model against which the known models are compared, thereby obtaining an absolute measure of model performance in a Bayesian framework.  We apply this new methodology to the problem of the dark energy equation of state, comparing an absolute upper bound on the Bayesian evidence for a presently unknown dark energy model against a collection of known models including a flat \\lcdm~scenario. We find a strong absolute upper bound to the Bayes factor $B$ between the unknown model and~\\lcdm, giving $B \\lsim 3$. The posterior probability for doubt is found to be less than 6\\% (with a 1\\% prior doubt) while the probability for \\lcdm~rises from an initial 25\\% to just over 50\\% in light of the data. We conclude that \\lcdm~remains a sufficient phenomenological description of currently available observations and that there is little statistical room for model improvement. ",
        "introduction": "One of the most important questions in cosmology is to identify the fundamental model underpinning the vast amount of observations nowadays available. The so-called ``cosmological concordance model'' is based on the cosmological principle (i.e., that the Universe is isotropic and homogeneous, at least on large enough scales) and on the hot Big Bang scenario, complemented by an inflationary epoch. This remarkably simple model is able to explain with only half a dozen free parameters observations spanning a huge range of time and length scales. Since both a cold dark matter (CDM) and a cosmological constant ($\\Lambda$) component are required to fit the data, the concordance model is often referred to as ``the $\\Lambda$CDM model''. It is however important to keep in mind that at this stage the \\lcdm~model is not a model in the sense attributed to the word by particle physicists, but rather a phenomenological scenario that appears to be able to explain the vast majority of observations with a great economy of free parameters.  In the classical approach to statistics, models (or hypotheses) can never be proved true, only falsified. \\cite{Popper}, for example, argued that theories always remain ``infinitely improbable'' regardless of the amount of evidence gathered in their favour. However, in the context of Bayesian inference support can be accrued for a model if the observed data verify predictions made by the model but not by competing models (see~\\cite{JaynesBook}). This is the subject of Bayesian model selection (see e.g.~\\cite{TrottaSky2008, Trotta2007} for applications to the cosmological context): given a set of competing models, the Bayes factor gives a measure of the relative performance of each model in explaining the data.  This program naturally prefers models that provide a good fit with the fewest number of free parameters,  thus implementing a quantitative version of Occam's razor.  Although Bayesian model selection can identify the best model in a given set of known models, it has no way of indicating whether the absolute quality of the preferred model is high or low. However, it seems desirable to be able to gauge the absolute performance of a model in a Bayesian sense, rather than just its relative performance with respect to  known competitors. In particular, this seems crucial for deciding whether the set of known  models includes the true model.  The purpose of this paper is to build on the notion of Bayesian doubt introduced by~\\cite{StarkmanTrotta2008} to develop and apply a Bayesian technique for model discovery, focusing in particular on the nature of dark energy. The structure of this paper is as follows: in section~\\ref{sec:doubt} we recall the notion of Bayesian doubt and introduce a new procedure for estimating an upper bound for the Bayes factor in favour of doubt. We next employ this procedure in section~\\ref{sec:cosmo} to assess the state of our knowledge of the dark energy equation of state,  focusing on the status of the current \\lcdm~concordance model. We present our results in terms of the posterior probability for doubt and for \\lcdm~in section~\\ref{sec:results} and discuss our conclusions in section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:results}  We now proceed to evaluate the doubt and the posterior probability of \\lcdm~for various combinations of cosmological data sets.  \\subsection{Model comparison outcome including doubt}   In Table~\\ref{tab:evid}, we present the estimated upper limit on the Bayes factor between \\lcdm~and model $\\mX$ as well as   the Bayes factors with respect to \\lcdm~for the other known models. Among the known models, we confirm what many others have shown -- that \\lcdm~is the best-fit known model, or at least that no other model is demonstrably better. Thus, we  find an inconclusive model comparison result (according to the Jeffreys' scale, Table~\\ref{tab:jeff}) when comparing \\lcdm~and a model with a free (but constant) $w$. We also find weak to moderate evidence  ($1 \\lsim \\ln B \\lsim 2.5$) against spatially curved models when compared to a flat \\lcdm, in agreement with the more detailed findings of~\\cite{Vardanyan:2009ft}.  Finally, there is weak to moderate evidence against the most complex of the known models -- one exhibiting both $w\\neq -1$ and $\\Omk \\neq 0$. This is in good agreement with the results of previous more thorough analyses, e.g.~\\cite{Liddle:2006tc,Liddle:2006kn} and ~\\cite{Li:2009}. From this, ordinary Bayesian model comparison concludes that \\lcdm~is still the best of the known models (at least for the limited range of alternative models considered here).  Most importantly, in the table, we report the improvement in the best-fit log-likelihood obtained over \\lcdm~by using $\\mX$, and use this to compute an absolute upper bound to the Bayes factor via Eq.~\\eqref{upperevi}. We notice that the improvement in the best-fit is fairly modest for all the data sets considered, supportive of the general sentiment in the community that \\lcdm~is in good agreement with available observations and that therefore there is little room for statistical improvement of the quality of fit.  This is in part because it is very hard to improve the quality of fit by changing $w(z)$ -- observables are usually a double integral of $w(z)$, and therefore insensitive to  features in the equation of state (see e.g~\\cite{Huterer:1999};~\\cite{Maor:2001} and \\cite{Clarkson:2009jq}). As a consequence, even a highly flexible $w(z)$ model such as the one we used here to describe $\\mX$ will lead to only small observable departures from the standard cosmological constant scenario.   It is important to keep in mind that such statements depend strongly on the statistics one employs to examine the models.  For example, the standard likelihood function for CMB data is insensitive to most of the reported anomalies in the low-$\\ell$ CMB \\cite{Copi:2010na,Bennett:2010jb}.  The interesting consequence from the point of view of doubt is that this translates into strong upper limits for the Bayes factor between model $\\mX$ and \\lcdm~(third from last column of Table~\\ref{tab:evid}). We find that the upper limit on the Bayes factor  $\\bBxl$ (last column of Table~\\ref{tab:evid}) for all the data combinations is less than $3$, just around ``weak evidence'' threshold ($\\ln B = 1$, see Table~\\ref{tab:jeff}). From our discussion in section~\\ref{sec:doubtdiss}, this means that the necessary condition for doubt to grow, $p(\\mX) \\bBxl \\gg 1$, is not met for any reasonable doubt prior choice.  We remind the reader at this point that our unknown model $\\mX$ has been designed in such a way as to exhibit the maximum possible evidence against \\lcdm. Therefore, if even such a model cannot achieve a significant level of evidence against \\lcdm, one can safely conclude that no other reasonable model will.  Of course this conclusion depends both on  the set of observations we have considered and on the particular likelihood function we have ascribed to that data.  New statistical treatments can bring to light anomalies in the existing data, while new observations might contain new unexpected features. \\begin{table*} % \\begin{tabular}{l | c c c | c c | c c } \\hline\\hline & $-1.3<w<-0.7$  & $w=-1.0$ & $-1.3<w<-0.7$ &  \\multicolumn{2}{c}{``Unknown'' model $X$} & $\\avB$ & $\\bBxl$\\\\ &  $\\Omega_{\\kappa}=0.0$ & $-0.3 < \\Omega_{\\kappa} <0.3$ & $-0.3 < \\Omega_{\\kappa} <0.3$  \\\\ &$\\ln B_{j\\Lambda}$    &$\\ln B_{j\\Lambda}$&$\\ln B_{j\\Lambda}$  &$\\Delta\\chi^2$  & $\\ln \\bBxl$   \\\\ \\hline CMB only       &$ 0.18 \\pm 0.09$      &$-1.03 \\pm 0.09$   &    $-1.09 \\pm 0.09$ &$ -1.21 \\pm 0.3$&$ 0.61  \\pm 0.2$ & $0.72 \\pm 0.03$ &$ 1.83 \\pm  0.3$\\\\ CMB + SN       &$-0.37 \\pm 0.09$      &$-1.30 \\pm 0.09$   &    $-1.63 \\pm 0.09$ &$ -2.34 \\pm 0.3$&$ 1.17  \\pm 0.2$ & $0.54 \\pm 0.02$ &$ 3.22 \\pm  0.5$\\\\ CMB + mpk      &$-0.50 \\pm 0.08$      &$-2.57 \\pm 0.08$   &    $-2.69 \\pm 0.08$ &$  -0.88\\pm 0.3$&$ 0.44  \\pm 0.2$ & $0.44 \\pm 0.01$ &$ 1.55 \\pm  0.2$\\\\ CMB + SN + mpk &$-0.48 \\pm 0.09$      &$-2.51 \\pm 0.09$   &    $-2.73 \\pm 0.09$ &$ -2.15 \\pm 0.3$&$ 1.08  \\pm 0.2$ & $0.44 \\pm 0.01$ &$ 2.93 \\pm  0.6$\\\\ \\hline \\end{tabular} \\caption{In the first three columns, we report the Bayes factors between the known models and \\lcdm~for different combinations of data sets, where $\\ln B_{j\\Lambda}<0$ favours \\lcdm. The fourth columns gives $\\Delta\\chi^2 = \\chi^2_\\mX - \\chi^2_\\Lambda$, the improvement in the best-fit log-likelihood obtained by using model $\\mX$ (specified in the text) over \\lcdm. The last column gives the corresponding absolute upper bound to the Bayes factor between model $\\mX$ and \\lcdm.} \\label{tab:evid} \\end{table*}   %   Our results in terms of posterior probability for doubt and for the \\lcdm~model are shown in Table~\\ref{tab:unknownflat}, for two different assumptions regarding the level of prior doubt, $p(\\mX) = 10^{-2}$ and $p(\\mX) = 10^{-6}$. These two choices are representative of a range that we think might bracket reasonable prior expectations: a prior doubt of 1\\% is certainly not too large, while leaving a little space for updating our models beliefs in the light of data. A prior doubt of $10^{-6}$ reflects the fact that surely we have to allow for a one-in-a-million chance that our current list of known models might be incomplete, and that the true underlying dark energy model might still be undiscovered.  Table~\\ref{tab:unknownflat} contains the level of doubt, which is updated from the prior by using the results of Table~\\ref{tab:evid} for the models' evidences. We find an increase in doubt by a factor of $\\sim 6$ for the most constraining data combination (CMB+SN+mpk). This however is largely a consequence of the doubt model acquiring some of the probability mass of the known models other than \\lcdm, as discussed under Case 3 in section~\\ref{sec:doubtdiss}. Indeed, the posterior probability of \\lcdm~is observed to increase (last column of Table~\\ref{tab:unknownflat}), from the initial prior value $p(\\Lambda) \\approx 0.25$ to just over 50\\% for the most constraining data combination. This result is almost independent of the choice of prior doubt. The behaviour of the posterior probability for doubt and \\lcdm~for a prior choice $p(\\mX) = 10^{-2}$ is shown in Fig.~\\ref{fig:doubt}, as a function of the data sets employed.  \\begin{table*} % \\begin{tabular}{lcccc} \\hline\\hline &\\multicolumn{2}{c} {Doubt $\\dt$} & Posterior for \\lcdm, $p(\\Lambda | d)$     \\\\ & Prior doubt: $p(\\mX)= 10^{-2}$& Prior doubt: $p(\\mX)=10^{-6}$ & (with $p(\\mX) = 10^{-2}$ and $p(\\Lambda) \\approx 0.25$) \\\\ \\hline CMB only  &$(2.50 \\pm0.2)\\times10^{-2}$&$(2.54\\pm0.2)\\times10^{-6}$ & $0.34 \\pm 0.01$\\\\ CMB+SN    &$(5.69 \\pm0.5)\\times10^{-2}$&$(5.97\\pm0.6)\\times10^{-6}$ & $0.44 \\pm 0.01$\\\\ CMB+mpk   &$(3.46 \\pm0.3)\\times10^{-2}$&$(3.54\\pm0.4)\\times10^{-6}$ & $0.55 \\pm 0.02$\\\\ CMB+SN+mpk&$(6.29 \\pm0.8)\\times10^{-2}$&$(6.64\\pm0.9)\\times10^{-6}$ & $0.53 \\pm 0.02$ \\\\       \\hline \\end{tabular} \\normalsize \\caption{First two columns: Posterior doubt for different data sets combinations and two prior doubt assumptions. Last column: posterior probability for the \\lcdm~model when allowing for the possibility of a 1\\% prior doubt on the completeness of our list of known models. } \\label{tab:unknownflat} \\end{table*} \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{posteriorplot} \\caption{Posterior probability for doubt for the \\lcdm~model as a function of different combinations of data sets. The probability of \\lcdm~increases from the initial 25\\% to just about over 50\\%, while the probability of doubt increases from the initial 1\\% to just over 6\\%, mostly as a consequence of acquiring probability from the other 3 known models considered in the analysis. This signals that \\lcdm~remains the most valid statistical description of the data.} \\label{fig:doubt} \\end{figure}  \\subsection{Impact of the addition of further known models} \\label{sec:moremodels}  \\begin{table} \\begin{tabular}{l l  l l} Known models    & \\multicolumn{3}{c}{Required $\\Delta\\chi^2$ for $p(\\mX | d) = p(\\Lambda | d)$} \\\\ $N$ & $p(\\mX) = 10^{-2}$ & $f$ & $p(\\mX) = p(\\Lambda)/f$  \\\\\\hline 4 & $-6.4$  & 4 & $-5.5$\\\\ 10 & $-4.6$ & 10 & $-9.2 $ \\\\ 20  & $-3.2$ & $10^2$ & $-18.4$\\\\ 50  & $-1.4$ & $10^3$ & $-27.6$ \\\\ \\hline \\end{tabular} \\caption{Improvement in the $\\chi^2$ of \\lcdm~required for the unknown model $\\mX$ to have the same {\\em a posteriori} probability as \\lcdm. First two columns: as a function of the number of known models, $N$, assuming a fixed prior doubt $p(\\mX) = 10^{-2}$. Last two columns: assuming a fixed fractional prior doubt, $p(\\mX) = p(\\Lambda)/f$, and as a function of $f$.  It is assumed that the evidence of the known models is much smaller than the evidence for \\lcdm. \\label{tab:chisq} } \\end{table}  We now proceed to estimate the robustness of our findings with respect to expanding the set of known models.  As has been mentioned above, the list of three alternative known models to \\lcdm~we adopted in this work is far from complete. However, even if a larger number of models $N$ were included in the known models list, it is reasonable to assume that the value of the average evidence between the known models and \\lcdm~would scale approximately as $\\propto 1/N$, for there is no other known model that presently can achieve a substantially higher evidence than \\lcdm~(if this was the case, then this other best model would take the place of \\lcdm~and become our baseline model which we seek to doubt -- or rather the dominant model in our list of models where we intend to compute the doubt for the whole list). By equating Eqs.~\\eqref{Doubt2} and \\eqref{eq:post_Lambda} we can solve for the value of $\\Delta\\chi^2$ required for the posterior on doubt to be equal to the posterior of \\lcdm. This gives the approximate condition (assuming that $\\avB \\approx 1/N$ and that $p(\\mX)  \\ll 1$): \\be \\Delta\\chi^2 \\approx 2 \\ln (p N). \\ee So the value of $\\Delta\\chi^2$ required for posterior doubt to reach the posterior for \\lcdm~scales logarithmically with the number of known models. Assuming a prior doubt $p(\\mX) = 10^{-2}$ one obtains the values of $\\Delta\\chi^2$ listed in the first column of Table~\\ref{tab:chisq} as a function of $N$. As more known doubts are put on the table, it becomes easier to doubt \\lcdm. From this scaling, it would appear that the improvement of $\\Delta \\chi^2 = -2.3$ for model $\\mX$ reported in Table~\\ref{tab:evid} for the data combination cmb+SN would lead to a larger probability of doubt than for \\lcdm~if we had assumed a list of $N \\gsim 30$ known models, rather than just three. As illustrated in Fig.~\\ref{fig:doubt_behaviour}, this effect is however a consequence of our choice of spreading the level of prior probabilities among the $N$ known models, while assuming {\\em a fixed} $p(\\mX)$, see Eq.~\\eqref{eq:priors}. As $N$ increases, the prior for \\lcdm~decreases while the prior doubt is kept constant. As a consequence, it becomes easier for the former to ``catch up'' with the latter. \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{postdoubtlcdm} \\caption{Posterior for doubt (dashed lines) and for \\lcdm~(solid lines) as a function of $-\\Delta\\chi^2 = \\chi^2_\\Lambda - \\chi^2_X$ assuming a fixed prior doubt $p(\\mX) = 10^{-2}$. Different curves are for different numbers of known models, $N=1,4,10,20$ (from thin to thick), assuming that $\\avB \\approx 1/N$.} \\label{fig:doubt_behaviour} \\end{figure}  In order to avoid this spurious effect, one could choose to set the prior doubt as a fraction $1/f$ ($f>1$) of the prior probability for \\lcdm, i.e., to require that the relative probability between $\\mX$ and $\\Lambda$ is constant {\\em a priori}, independent of the number of known models. We thus replace the prescription of Eq.~\\eqref{eq:priors} by \\begin{align} p(\\Lambda) & = \\frac{1}{N}(1-p(\\mX)) \\\\ p(\\mX) & = \\frac{p(\\Lambda)}{f} = (Nf + 1)^{-1} \\end{align} and by equating the posterior doubt with the posterior for \\lcdm~we obtain the following requirement for the $\\Delta\\chi^2$: \\be \\Delta \\chi^2 = -4 \\ln f. \\ee This is now independent of the number of known models $N$ and it only depends logarithmically on the prior doubt fraction, $f$. From the last two columns of Table~\\ref{tab:chisq} we can see that even if doubt started off a factor of just $f=4$ less probable than \\lcdm, a $\\Delta\\chi^2 = -5.5$ would be required in order for the unknown model $\\mX$ to become as probable as \\lcdm. Increasing the prior gap between doubt and \\lcdm~(i.e., increasing $f$) only makes the requirements on the $\\chi^2$ improvement more taxing.  In summary, once the effect of adding extra models to the known models' list is corrected for by introducing the fractional prior doubt $f$, we find that the improvement in the $\\chi^2$ found for various combinations of data sets is insufficient to doubt \\lcdm. If the unknown model starts off being a factor of 4 less probable than \\lcdm, one would need an improvement in the $\\chi^2$ of about 5 units to reverse the situation in the posterior, which is quite a bit larger than the maximum $\\chi^2$ improvement observed from the data."
    },
    "1005/1005.4419_arXiv.txt": {
        "abstract": "Galactic outflows of low ionization, cool ($\\sim10^4\\K$) gas are ubiquitous in local starburst galaxies, and in the majority of galaxies at high redshift. How these cool outflows arise is still in question. Hot gas from supernovae has long been suspected as the primary driver, but this mechanism suffers from its tendency to destroy the cool gas as the latter  is accelerated. We propose a modification of the supernova scenario that overcomes this difficulty.  Star formation is observed to take place in clusters; in a given galaxy, the bulk of the star formation is found in the $\\sim 20$ most massive clusters. We show that, for $L^\\star$ galaxies, the radiation pressure from clusters with $\\mcl \\gtrsim 10^6\\,M_\\odot$ is able to expel the surrounding gas at velocities in excess of the circular velocity of the disk galaxy. This cool gas can travel above the galactic disk in less than $2\\Myr$, well before any supernovae erupt in the driving cluster. Once above the disk, the cool outflowing gas is exposed to radiation, and supernovae induced hot gas outflows, from other clusters in the disk, which in combination drive it to distances of several tens to hundreds of kiloparsecs. Because the radiatively driven clouds grow in size as they travel, and because the hot gas is more dilute at large distance, the clouds are less subject to destruction if they do eventually encounter hot gas. Therefore, unlike wind driven clouds, radiatively driven clouds can survive to distances $\\sim 50\\kpc$, and potentially give rise to the metal absorbers seen in quasar spectra. We identify these cluster-driven winds with large-scale galactic outflows. In addition, we note that the maximum cluster mass in a galaxy is an increasing function of the galaxy's gas surface density. As a result, only starburst galaxies, where massive clusters reside, are able to drive winds cold outflows on galactic scales via this mechanism. We find that the critical star formation rates above which large scale cool outflows will be launched to be $\\dot\\Sigma^{crit}_*\\approx 0.1\\, M_\\odot\\yr^{-1}\\kpc^{-2}$, which is in good agreement with observations. ",
        "introduction": "Observational studies of both nearby and high redshift star forming galaxies have established that cold gas emerges from such galaxies at velocities ranging from a few tens to several hundred kilometers per second \\citep{1990ApJS...74..833H,1996ApJ...462L..17S,1997ApJ...486L..75F,2000ApJS..129..493H,2000ApJ...528...96P,2001ApJ...554..981P,2003ApJ...588...65S,2005ApJ...621..227M,2005ApJS..160..115R,2007ApJ...663L..77T,2009ApJ...692..187W,2009arXiv0912.3263M}. Hot outflowing gas from supernovae (SN) has been considered to be the primary mechanism for driving galactic outflows (\\cite{1985Natur.317...44C,2002ASPC..254..292H} and references therein). It is generally assumed that cold gas and dust are entrained in the hot gas.  As numerous authors have pointed out, in such a scenario the cold gas is subject to Kelvin-Helmholtz instabilities \\citep{1994ApJ...420..213K,2008ApJ...674..157C} and/or by conductive evaporation \\citep{2005MNRAS.362..626M} and ends up surviving for less than $1\\Myr$ (a few cloud crushing times) before being destroyed. The cloud crushing time follows $\\tcc=(\\rho_c/\\rho_h)^{1/2}\\,\\rc/v_b$, where $\\rho_c/\\rho_h$ is the density ratio between the cold gas and the hot gas, $\\rc$ is the radius of the cloud before it is exposed to the hot gas, and $v_b\\approx1000\\kms$ is the shock velocity in the hot gas. This implies that the cold gas can travel less than a few hundred parsecs before being mixed into the hot gas and hence effectively destroyed. However, observational evidence shows that cold gas survives at least out to ten $\\kpc$. An explanation for the presence of this cold gas is required. In this paper we show how taking into account the effect of radiation pressure on dust grains can affect the fate of the gas.  A second point regarding the driving of outflows is that most stars form in massive clusters, both in quiescent spirals that lack strong outflows, like the Milky Way, and in starburst galaxies such as M82. We argue that outflows emerge from the most massive clusters in a galaxy, as opposed to arising from the collective effects of smaller clusters. The radiation pressure from small clusters will not punch holes in the gas disk, and the subsequent supernovae will either leak into the interstellar medium (ISM) in low star formation rate galaxies, or will  radiate their energy away after their growth has been halted by the high pressure of the ISM in high star formation rate galaxies (such as Arp 220). In contrast, the radiation pressure from massive clusters will blow cold gas out of the disk, paving the way for the hot gas from SN to escape.   In this paper we present a model including both radiation pressure from massive clusters and ram pressure from supernovae. As recently pointed out by \\cite{2009MNRAS.396L..90N}, we find that before any cluster stars explode as SNe, radiation pressure, acting alone, launches the cold gas surrounding a star cluster to heights of several hundred parsecs above the disk. As the cold cloud travels above the disk, it expands. This expansion, as well as a lower hot gas density at larger radii, significantly increase the cloud crushing time once it does encounter hot gas, allowing it to survive long enough to reach distances exceeding one hundred $\\kpc$. Once above the disk, a cold gas cloud experiences a ram pressure force which is comparable to the radiation pressure force \\citep{2005ApJ...618..569M}.  While it is clear that both driving mechanisms are significant, the action of the radiation pressure is crucial for the survival of cold gas out to large radii. We also illustrate the importance of this force by showing that it is capable of driving outflows even in the absence of hot gas. In \\S \\ref{sec:starclusters} we describe how the masses of the largest star clusters in a galaxy are determined by the interplay between self-gravity and radiation pressure in giant molecular clouds. In \\S \\ref{sec:winds} we describe how winds are launched from massive clusters, and how this leads to a critical star formation rate to produce galactic winds. We follow the evolution of these winds outside the star forming disk, to distance of order the galactic virial radius, in \\S \\ref{sec:galacticwinds}. We describe the cloud properties, cover factor, and mass loss rates of the wind in \\S \\ref{sec:properties}. We discuss our results in \\S \\ref{sec:discussion} and offer conclusions in the final section. ",
        "conclusions": "\\label{sec:conclusions}  We have shown that radiation pressure from massive stellar clusters can drive outflows with velocities $v\\sim{\\rm a~few}\\times 100\\,\\kms$ to distances $\\sim50-100\\kpc$ from their host galaxies. The outflows consist of cold ($\\sim10^4\\,\\K$) and dusty clouds with characteristic column densities in the range $N_H\\sim10^{21}\\cm^{-2}$ near the galactic disk, ranging down to $N_H\\sim10^{18}\\cm^{-2}$ at $r\\sim 50\\kpc$. The gas in the clouds is in a low ionization state. Once a cloud is above the disk of a galaxy, it experiences the radiation pressure from nearby clusters as well as ram pressure from supernovae.  This solves a major problem of hot outflow-driven cold clouds, namely their short lifetimes. Since radiatively driven clouds need not be in contact with (relatively) dense hot gas in order to be driven out of the galaxy, they are not destroyed by Kelvin-Helmholtz and/or conduction.  We have argued that individual clusters can drive outflows above the disk. If the disk is sufficiently luminous, the disk light can then drive the outflow to distances of tens of kiloparsecs. This evades the need for a global ``blowout'' of the interstellar medium of (the central portions of) the galaxy in order to drive a large scale wind. We note that ULIRG are clearly not ``blown-out'' and yet have winds, so this is an attractive feature of cluster driven winds.  Since the gas is launched from individual star clusters roughly vertically from the disk, it will retain the rotational motion of the disk. If the launching cluster is at large disk radius, measurements of the rotational velocity will find gas at large radii and large azimuthal velocities. The mass loss rates from cluster driven winds are comparable to the star formation rate in massive galaxies, and substantially larger than the star formation rate in low mass galaxies. In both types of objects, the terminal velocity of the wind is similar to the rotational velocity of the galaxy. While the winds may or may not reach the escape velocity, they can nevertheless reach radii in excess of the host galaxy's virial radius.  Since the cluster escape velocity must be comparable to the escape velocity from the galaxy to launch a flow above the disk, and because the maximum cluster mass is related to the gas surface density and hence star formation rate, there is a critical star formation rate (per unit area) $\\dot \\Sigma_*\\approx0.1\\,M_\\odot\\yr^{-1}\\kpc^{-2}$ to launch a wind radiatively. This is in good agreement with the threshold inferred from observations.  This paper has employed simply analytic and one dimensional hydrodynamic models of cluster driven winds to argue that such winds are crucial to the formation of galactic superwinds. The results provide strong motivation to carry out more realistic three dimensional radiative hydrodynamic modeling."
    },
    "1005/1005.0710_arXiv.txt": {
        "abstract": "{Hydrodynamical simulations of two giant planets embedded in a gaseous disk have shown that in case of a smooth convergent migration they end up trapped into a mean motion resonance. These findings have led to the conviction that the onset of dynamical instability causing close encounters between the planets can occur only after the dissipation of the gas when the eccentricity damping is over.} { We show that a system of three giant planets may undergo planet-planet scattering when the gaseous disk, with density values comparable to that of the Minimum Mass Solar Nebula, is still interacting with the planets.} {The hydrodynamical code FARGO--2D--1D is used to model the evolution of the disk and planets, modified to properly handle close encounters between the massive bodies.} {Our simulations predict a variety of different outcomes of the scattering phase which includes orbital exchange, planet merging and scattering of a planet in a hyperbolic orbit.} {This implies that the final fate of a multiplanet system under the action of the disk torques is not necessarily a packed resonant configuration.} ",
        "introduction": "\\begin{figure*}[hp] \\begin{center} \\begin{tabular}{c c} \\hskip -2 truecm \\resizebox{80mm}{!}{\\includegraphics[angle=-90]{14716fig01.eps}} % \\resizebox{80mm}{!}{\\includegraphics[angle=-90]{14716fig02.eps}} \\\\ \\hskip -2 truecm \\resizebox{80mm}{!}{\\includegraphics[angle=-90]{14716fig03.eps}} % \\resizebox{80mm}{!}{\\includegraphics[angle=-90]{14716fig04.eps}} \\\\ \\hskip -2 truecm \\resizebox{80mm}{!}{\\includegraphics[angle=-90]{14716fig05.eps}} % \\resizebox{80mm}{!}{\\includegraphics[angle=-90]{14716fig06.eps}} \\\\ \\end{tabular} \\caption[]{Orbital evolution of the three planets while embedded in the disk. The top plots show a model where $\\Sigma_0 = 0.4 \\Sigma_{\\rm MMSN}$ and $a_1 = 3.0{\\rm AU}, a_2 = 6.7{\\rm AU}, a_3 = 10.5{\\rm AU}$ The middle plots show a case with $\\Sigma_0 = 0.65 \\Sigma_{\\rm MMSN}$ and $a_1 = 2.0{\\rm AU}, a_2 = 6.5{\\rm AU}, a_3 = 10.0{\\rm AU}$ The bottom plots illustrate the evolution when $\\Sigma_0 = \\Sigma_{\\rm MMSN}$ and $a_1 = 3.0{\\rm AU}, a_2 = 5.5{\\rm AU}, a_3 = 10.0{\\rm AU}$. } \\label{f1} \\end{center} \\end{figure*}  A multi--planet system made of Jupiter-sized bodies may reach completion when the planets are still under the influence of the disk.  In this scenario the mutual gravitational perturbations compete with the gravitational torques exerted on the planets by the disk. These torques  cause planet migration which is a strong source of mobility to the system and possibly causes the planets to come closer to each other. The mutual approach is supposed to end  when a pair of planets is captured in a mean motion resonance (e.g. the 5:3, 3:2, 3:1 or 2:1). The most extensively studied scenario is that of two jovian mass planets embedded in a viscously evolving disk.  \\cite{kpb, klm}, \\cite{ms} and \\cite{pn} have shown via self--consistent disk--planet hydrodynamical simulations that the two planets become trapped into a resonance and create a cavity encircling their orbits. This resonant configuration appears stable in these simulations which argues in favor of a scenario where the resonance locking should be broken only after gas dissipation possibly by the end of the gas damping or by planetesimal scattering.  Damped N--body models have also been used to study the evolution of two planet systems where the disk--planet interaction is modeled as an external dissipative force. \\cite{lp}, \\cite{bmf} and \\cite{ltr} have shown that the resonance trapping is permanent if a strong damping of eccentricity is adopted. Even turbulence appears not to prevent the resonance capture \\citep{rp} in this scenario. The escape from resonance can occur only if the eccentricity damping term used in the model is relatively low and, after the resonance breaking, planets may have close encounters. This second approach characterizes the outcome of the simulations by \\cite{ma}, where two planets may undergo planet--planet scattering after the escape from resonance trapping. However, hydrodynamical simulations have so far never shown such a behavior with the planets assumed to remain in resonance until the dissipation of the gas. The gas damping appears to be dominating over the mutual gravitational perturbations even for Jupiter size planets, which supports the idea that the preferred outcome of the evolution of a multiplanet system under the tidal action of the disk is a packed resonant system.  A system of three giant planets has proved to be dynamically more prone to instability and to develop crossing orbits \\citep{wm,cfmr} compared to the case of two planets. The mutual gravitational perturbations are stronger and may more easily overcome the disk--damping effect. The observed systems Ups And, 55 Cnc, HIP 14810, HD74156, HD160691 and HD 37124 also show that it is indeed possible to form more than two giant planets around solar type stars. For this reason we selected a scenario where three giant planets with masses close to that  of Jupiter  form in the disk and evolve under planet-disk and planet-planet interactions. We will show that the chaotic evolution, which is characterized by mutual close encounters between (giant) planets, is favored by planet migration. Our hydrodynamical simulations show that most of the time the planets are quickly captured in a three-body resonance during their migration,  despite their rapid drift rate, but this may become unstable. The growth in eccentricity leads the planets into crossing orbits, and they begin to evolve under the effects of mutual approaches.  The following chaotic phase ends up either with the ejection of a planet, planet--merging or orbital reconfiguration. ",
        "conclusions": "Although there are known two-planet systems in simple mean-motion resonances, there are also many systems that are either multiple and non-resonant, or that display evidence for scattering that would be suppressed if resonant capture was a ubiquitous process. We show that resonant capture is more easily evaded in a three-planet as opposed to a two-planet system, leading to a phase of planet--planet scattering. A major implication of planet--planet scattering in the gas disk is the significant re--arrangement of planet orbits and eccentricity excitation, which may affect the subsequent evolution of the planets.  As a consequence, the final mass distribution of planets as a function of the semi-major axis cannot be representative of the formation process. In addition, a closely packed planetary system where the multiple planets are all locked in mean--motion resonances is not the preferred outcome of planet migration, as previously argued. The chaotic phase may lead to different orbital configurations where the planets can be well separated into semi-major axis and stable even after the gas dissipation. Lower eccentricities are expected as the final state of the planet evolution because damping can occur after the  dynamically active phase \\citep{moe}. However, it may not be the case for corotation resonance saturation. According to  \\cite{gs, ol} a planet may experience a net growth of eccentricity via disk--planet interaction if its eccentricity is not small to begin with. This would be the case for a planet at the end of the chaotic phase.  Additional simulations are undergoing in order to perform a deeper exploration of the parameter space to identify the initial configurations of the planets and disk that lead to planet--planet scattering rather than to a stable resonant trapping."
    },
    "1005/1005.2009_arXiv.txt": {
        "abstract": "We report the discovery of \\hatcurb{}, a transiting extrasolar planet orbiting the $V=\\hatcurCCmag\\ \\rm{mag}$ \\hatcurISOspec\\ dwarf \\hatcurCCgsc, with a period $P=\\hatcurLCP\\,\\rm{d}$, transit epoch $T_c =\\hatcurLCT$ (BJD\\altaffilmark{$\\star\\star\\star$}), and transit duration \\hatcurLCdur\\,d. The host star has a mass of \\hatcurISOm\\,\\msun, radius of \\hatcurISOr\\,\\rsun, effective temperature \\hatcurSMEteff\\,K, and metallicity $\\feh= \\hatcurSMEzfeh$. The planetary companion has a mass of \\hatcurPPmlong\\,\\mjup, and radius of \\hatcurPPrlong\\,\\rjup\\ yielding a mean density of \\hatcurPPrho\\,\\gcmc. Comparing these observed characteristics with recent theoretical models, we find that \\hatcurb\\ is consistent with a 1 Gyr H/He-dominated gas giant planet. \\hatcurb\\ resides in a sparsely populated region of the mass--radius diagram and has a non-zero eccentricity of $e=0.036$ with a significance of $10\\sigma$. ",
        "introduction": "\\label{sec:introduction}  Planets that transit their host stars play a special role in our understanding of the characteristics of exoplanets: their transit allows us to accurately determine the radius and the orbital inclination of the planet from the photometric light curve so that an actual mass can be derived from a spectroscopic orbit of the host star. The mass and radius enable us to infer a bulk composition of the planet, and although there are degeneracies associated with the bulk composition, it allows us to put constraints on models of planetary structure and formation theories. The incredible diversity of the over 60 discovered transiting planets, ranging from dense planets with a higher mean density than that of copper to strongly irradiated puffed-up planets with a mean density comparable to that of corkwood, have baffled the exoplanet community, and no unified theory has been established to explain all the systems consistently. Transiting extrasolar planet (TEP) discoveries are primarily the result of dedicated ground--based searches, such as SuperWASP \\citep{pollacco:2006}, HATNet \\citep{bakos:2004}, TrES \\citep{alonso:2004} and XO \\citep{pmcc:2005}, and space-borne searches, such as CoRoT \\citep{baglin:2006} and Kepler \\citep{borucki:2010}.  Since its commissioning in 2003, the Hungarian-made Automated Telescope Network \\citep[HATNet;][]{bakos:2004} survey has been one of the major contributors to the discoveries of TEPs. HATNet has discovered over a dozen TEPs since 2006 by surveying bright stars ($8\\lesssim I \\lesssim 12.5$\\,mag) in the Northern hemisphere and has now covered approximately 11\\% of the Northern sky. HATNet consists of six wide field automated telescopes; four of these are located at the Fred Lawrence Whipple Observatory (FLWO) in Arizona, and two on the roof of the Submillimeter Array hangar (SMA) of the Smithsonian Astrophysical Observatory (SAO) in Hawaii. In this paper we report a new TEP discovery of HATNet, called \\hatcurb.  The structure of the paper is as follows. In \\refsecl{obs} we summarize the observations, including the photometric detection, and follow-up observations. In \\refsecl{analysis} we describe the analysis of the data, such as the stellar parameter determination (\\refsecl{stelparam}), ruling out blend scenarios (\\refsecl{blend}), and global modeling of the data (\\refsecl{globmod}). We discuss our findings in \\refsecl{discussion}. ",
        "conclusions": "\\label{sec:discussion}  We present the discovery of \\hatcurb\\ with a period of $P= \\hatcurLCPshort\\ \\rm{d}$, a mass of $\\mpl=\\hatcurPPm\\ \\mjup$ and a radius of $\\rpl=\\hatcurPPr\\ \\rjup$. Currently, there are only a handful of planets residing near \\hatcurb\\ in the mass--radius diagram. CoRoT-2b \\citep{alonso:2008} and CoRoT-6b \\citep{fridlund:2010} have similar masses and periods, but both have eccentricities consistent with zero. HD80606 \\citep{naef:2001} and HD17156 \\citep{fischer:2007} are also similar in mass, but both planets have long periods and very eccentric orbits. WASP-10b ($P=3.09$\\,d, $e=0.06, \\mpl=2.96\\ \\mjup, \\rpl=1.08\\ \\rjup$, see \\citealt{christian:2009,johnson:2009}) is probably the transiting planet which most resembles \\hatcurb\\, with a similar period and eccentricity, albeit slightly lower mass and radius. At $V= \\hatcurCCmag\\ \\rm{mag}$, \\hatcur\\ is one of the brightest stars known to host a $\\sim 4\\ \\mjup$ planet.  We have compared \\hatcurb\\ to the theoretical models from \\citet{fortney:2007} by interpolating the models to the solar-equivalent semi--major axis of $a_{equiv}=\\hatcurPPaequiv\\ \\rm{AU}$, the result of which can be seen overplotted in \\reffigl{exomr}. We find that the mass and radius of \\hatcurb\\ are quite consistent with the 1 Gyr model and conclude that \\hatcurb\\ is most likely a H/He-dominated planet.  The radial velocity dataset used to characterize \\hatcurb\\ consist of three different velocity sets from two different telescopes (FIES medium and high resolution spectra from the NOT and HIRES spectra from Keck-I), requiring additional free parameters in the global fit to account for the arbitrary velocity offset between the measurements. Nevertheless, the different observations yield very similar semi-amplitude and we find an rms from the best fit model of only $\\hatcurRVfitrms$\\,\\ms.  The rich dataset consisting of 27 RV observations allows us to accurately determine the small non-zero eccentricity of the orbit as $e=0.036$ with a high significance of $10\\sigma$, with $k \\neq 0$ at $10\\sigma$ and $h \\neq 0$ at $3\\sigma$. The eccentricity is consistent and clearly evident even when fitting separate orbits to the three different velocity datasets.  Most short--period TEPs have eccentricities consistent with zero, as is expected because of circularization driven by tidal dissipation within the planet with a typical timescale substantially less than 1 Gyr (see \\citealt{jackson:2008, mardling:2007}). In fact, the estimated circularization timescale for \\hatcurb\\ is $\\tau_{cir} \\sim 0.03\\ \\rm{Gyr}$ (when the tidal dissipation parameter is assumed to be $Q_p = 10^5$), which is much less than the estimated $\\sim 2\\ \\rm{Gyr}$ age of the system. On the other hand, \\citet{matsumura:2008} argue that circularization timescales could be significantly higher and the planets with eccentric orbits may simply be in the process of being circularized. \\citet{adams:2006} argue that for multiple systems with a hot Jupiter as the inner planet, secular excitation of the eccentricity by companion planets could explain the non--zero eccentricity of these systems. Thus the origin of the eccentricity of \\hatcurb\\ remains unclear.  With a stellar rotation of \\hatcurSMEiivsin\\ \\kms and an impact parameter of $b \\sim 0.43$, \\hatcur\\ is a favorable target for measuring the Rossiter--McLaughlin (RM) effect. This effect can be used to determine the alignment of the planet's orbital angular momentum vector with the stellar spin axis \\citep{winn:2009}.  The estimated amplitude of the RM effect is $47\\ \\ms$ for \\hatcurb.  \\begin{figure*}[!ht] \\plotone{MRdia_HAT-P-16b_v2.EPS} \\caption{ Mass--radius diagram of currently known TEPs. HATNet planets are show in green and \\hatcurb\\ is labeled explicitly. The Solar system planets are shown as filled gray squares. Isodensity curves (in $\\rm{g\\ cm^{-3}}$) are plotted as dotted lines. Overlaid are planetary 1.0 Gyr (brown, the upper set of  lines) and 4.5 Gyr (green, the lower set of lines) isochrones from \\citet{fortney:2007} for H/He-dominated planets with core mases of $M_{c}=0\\ \\mearth$ (solid) and $M_{c}=10\\ \\mearth$ (dashed), interpolated to the solar equivalent semimajor axis of \\hatcurb. \\label{fig:exomr}} \\end{figure*}"
    },
    "1005/1005.2379_arXiv.txt": {
        "abstract": "Recent blazar detections by HESS, MAGIC, and VERITAS suggest that very-high-energy (VHE, $E > 100$~GeV) $\\gamma$-rays may be produced in most, if not all, types of blazars, including those that possess intense circumnuclear radiation fields. In this paper, we investigate the interaction of nuclear VHE $\\gamma$-rays with the circumnuclear radiation fields through $\\gamma\\gamma$ absorption and pair production, and the subsequent Compton-supported pair cascades. We have developed a Monte-Carlo code to follow the spatial development of the cascade in full 3-dimensional geometry, and calculate the radiative output due to the cascade as a function of viewing angle with respect to the primary VHE $\\gamma$-ray beam (presumably the jet axis of the blazar). We show that even for relatively weak magnetic fields, the cascades can be efficiently isotropized, leading to substantial off-axis cascade emission peaking in the {\\it Fermi} energy range at detectable levels for nearby radio galaxies. We demonstrate that this scenario can explain the {\\it Fermi} flux and spectrum of the radio galaxy NGC~1275. ",
        "introduction": "Blazars, a class of active galactic nuclei (AGNs) comprised of Flat-Spectrum Radio Quasars (FSRQs) and BL~Lac objects, exhibit some of the most violent high-energy phenomena observed in AGNs to date. Their spectral energy distributions (SEDs) are characterized by non-thermal continuum spectra with a broad low-frequency component in the radio -- UV or X-ray frequency range and a high-frequency component from X-rays to $\\gamma$-rays. They show rapid variability across the electromagnetic spectrum. In extreme cases, the very-high-energy (VHE) $\\gamma$-ray emission of blazars has been observed to vary on time scales of just a few minutes \\citep{albert07a,aharonian07}.  Leptonic and hadronic models \\citep[for a recent review see, e.g.][]{boettcher07} can generally successfully account for the overall SEDs observed from the known VHE $\\gamma$-ray blazars, which are almost all BL Lac objects. BL Lac objects span a wide range of synchrotron peak frequencies, from IR to X-rays. According to the location of the synchrotron peak, they are classified as Low frequency peaked BL Lacs (LBLs: synchrotron peak in the IR), intermediate BL Lac objects (IBLs: synchrotron peak in the optical/UV) or High frequency peaked BL Lacs (HBLs: synchrotron peak in the X-rays). Until very recently, all known VHE $\\gamma$-ray blazars were HBLs. However, the recent detections of the IBLs W~Comae, 3C~66A, and PKS~1424+240 by VERITAS \\citep{acciari08,acciari09,acciari10}, the LBLs BL~Lacertae and S5~0716+714 by MAGIC \\citep{albert07b,anderhub09}, and even the Flat-Spectrum Radio Quasars (FSRQs) 3C~279 by MAGIC \\citep{albert08} and PKS~1510-089 by HESS \\citep{wagner10} suggest that most blazars are intrinsically emitters of VHE $\\gamma$-rays.  The fact that most LBLs and FSRQs have not been detected as VHE $\\gamma$-ray sources might be primarily a result of the absorption of VHE $\\gamma$-rays by lower-frequency (IR -- optical -- UV) radiation. As those objects tend to be located at greater cosmological distances ($z \\gtrsim 0.2$), $\\gamma\\gamma$ absorption on the Extragalactic Background Light (EBL) becomes substantial, suppressing any intrinsically produced $> 100$~GeV emission \\citep[e.g.][]{dk04,stecker08,franceschini08,finke09,finke10}. However, possibly even more importantly, multi-GeV $\\gamma$-rays produced in the high-radiation-density environment within the broad line region (BLR) and the dust torus of a quasar, is expected to be strongly attenuated by $\\gamma\\gamma$ pair production \\citep[e.g.][]{pb97,donea03,reimer07,liu08,sb08}. The studies cited above show that for typical parameters expected in 3C279, photons above $\\sim 100$~GeV are expected to be strongly attenuated unless the $\\gamma$-ray emission region is located beyond the broad line region. If VHE $\\gamma$-rays are produced very close to the central engine, also $\\gamma\\gamma$ absorption by the direct accretion disk emission may become substantial \\citep{sb10}.  Given the now substantial number of non-HBL VHE $\\gamma$-ray blazars, it seems plausible that most LBLs and FSRQs and their misaligned parent population, radio galaxies, are producing VHE $\\gamma$-ray emission within their blazar zone. This view is also supported by the detection of VHE $\\gamma$-ray emission from two non-blazar AGNs, namely the radio galaxies M87 \\citep{Aharonian04} and Cen~A \\citep{Aharonian09}. The non-detection of VHE $\\gamma$-rays from most LBLs and FSRQs might then be due to the combined effects of local and intergalactic $\\gamma\\gamma$ absorption. As a consequence of $\\gamma\\gamma$ absorption in the local radiation field of the AGN, GeV -- TeV electron-positron pairs are injected into the AGN environment, which, in the dense radiation field within the BLR and the dust torus of quasars, may initiate inverse-Compton supported pair cascades.  The development of pair cascades induced by VHE $\\gamma$-ray emission from blazars has so far concentrated on the development of Mpc-scale pair halos resulting from the interaction of ($\\gtrsim 100$~TeV) $\\gamma$-rays with the Cosmic Microwave background \\citep[e.g.][]{aharonian94}, or of VHE ($\\gtrsim 100$~GeV) $\\gamma$-rays by the EBL \\citep{venters10}. Due to the long Compton cooling timescale of pairs on the CMB or the EBL, those authors could reasonably consider the produced cascades isotropic, leading to an extended pair halo around the AGN. Small-angle deflection of secondaries in weak intergalactic magnetic fields has been included in those considerations by \\cite{plaga95} and \\cite{elyiv09}. The development of pair cascades within the AGN has been discussed by \\cite{bk95} and \\cite{sb10}.  Depending on the magnetic field strength and orientation in the extended nuclear region, cascades developing within the high-radiation-energy-density environment within the BLR of a quasar may be efficiently isotropized in the immediate vicinity of the AGN. This will lead to distinct spectral features, which we will consider in this paper. In Section \\ref{estimates} we present some general considerations and analytic estimates of the expected results, including an estimate of the required magnetic field strengths within $\\sim$~a few pc from the central engine of a blazar (or radio galaxy), for which secondary electrons and positrons in pair cascades initiated by primary VHE $\\gamma$-rays may be efficiently isotropized within the central region. The consequence would be quasi-isotropic radiation signatures from Compton-supported pair cascades, potentially observable in the {\\it Fermi} energy range. In Section \\ref{setup} we will describe a numerical code that treats the full three-dimensional development of these cascades, together with the general model setup and simplifying assumptions. Numerical results for generic parameters will be presented in Section \\ref{parameterstudy}. In Section \\ref{NGC}, we will demonstrate that the recent {\\it Fermi} detection of the radio galaxy NGC~1275 \\citep{abdo09a} can be plausibly explained by the cascade emission from a misaligned VHE $\\gamma$-ray emitting blazar. We summarize and present an outlook towards future work in Section \\ref{summary}. ",
        "conclusions": ""
    },
    "1005/1005.5073_arXiv.txt": {
        "abstract": "{{SPIRE, the Spectral and Photometric Imaging Receiver, is the \\it{Herschel}\\rm~ {Space Observatory`s submillimetre camera and spectrometer.  It contains a three-band imaging photometer operating at 250, 350 and 500 $\\mu$m, and an imaging Fourier transform spectrometer (FTS) covering 194-671 $\\mu$m (447-1550 GHz). In this paper we describe the initial  approach taken to the absolute calibration of the SPIRE instrument using a combination of the emission from the \\it{Herschel}\\rm~ {telescope itself and the modelled continuum emission from solar system objects and other astronomical targets.  We present the photometric, spectroscopic and spatial accuracy that is obtainable in data processed through the ``standard'' pipelines. The overall photometric accuracy at this stage of the mission is estimated as 15\\% for the photometer and between 15 and 50\\% for the spectrometer.  However, there remain issues with the photometric accuracy of the spectra of low flux sources in the longest wavelength part of the SPIRE spectrometer band.  The spectrometer wavelength accuracy is determined to be better than 1/10$^{th}$ of the line FWHM. The astrometric accuracy in SPIRE maps is found to be 2 arcsec when the latest calibration data are used. The photometric calibration of the SPIRE instrument is currently determined by a combination of uncertainties in the model spectra of the astronomical standards and the data processing methods employed for map and spectrum calibration.  Improvements in processing techniques and a better understanding of the instrument performance will lead to the final calibration accuracy of SPIRE being determined only by uncertainties in the models of astronomical standards.}}}} ",
        "introduction": "The in-flight calibration of any astronomical instrument relies on a combination of previous measurements using instrumentation with known or traceable calibration accuracy, accurate models of standard astronomical sources and calibration sources internal to the instrument or facility. In this paper we discuss the methods and source models employed to convert SPIRE data to physical units and discuss the accuracy of the calibration and any caveats that must be placed on the data.  The SPIRE instrument and its overall calibration scheme are described in detail in Griffin et al. (\\cite{griffin10}).  For the purposes of discussion of the calibration and performance estimation we treat SPIRE as two separate instruments that are spatially separated at the focal plane of the \\it{Herschel}\\rm~ telescope: a three band imaging photometric camera with bands centred nominally on 250 {$\\mu$}m\\ (PSW), 350 {$\\mu$}m\\ (PMW), and 500 {$\\mu$}m\\ (PLW) with detector arrays of 139, 88 and 43 respectively feedhorn coupled NTD-bolometers (Turner et al \\cite{turner01}), and a two band imaging Fourier transform spectrometer (FTS) covering 194-313$\\mu$m\\ (SSW) and 303-671$\\mu$m\\ (SLW) with arrays of 37 and 19 detectors. SPIRE contains two internal calibration sources which were designed to provide a rapidly modulated signal for relative gain calibration (the PCAL source) and a stable thermal source for calibration of the spectrometer and to balance the power from the telescope (the SCAL source). PCAL is placed at an image of the telescope secondary in the common light path of the two instruments and therefore can be used to stimulate all detectors in the photometer and spectrometer. The SCAL source is placed in the second input port of the spectrometer (Swinyard et al. \\cite{swinyard03}). It is therefore constantly viewed during observations and provides a reference against which the spectrum from the telescope port is measured. The intention before flight was that the SCAL source would be heated to a temperature sufficient to entirely ``null'' the spectrum from the telescope.  We discuss in Sect. \\ref{spectro_pipe} why this has proved to be unnecessary for standard observing conditions.  The internal calibrators are only viewed through part of the optics chain and the overall astronomical calibration and performance of the instrument and telescope together can only be established by observation of astronomical sources as we describe in this paper.  The structure of the paper is as follows: in Sect. \\ref{initial} we describe the initial in-flight performance evaluation of the instrument before going on to cover the basic conversion of the data from digital encoder values to physical units for the detectors, the beam steering mirror (BSM) and the spectrometer mechanism (SMEC) in Sect. \\ref{dataproc}.  In Sect. \\ref{time_constant} we describe the characterisation of the time response of the photometer bolometers using ionising radiation hits and in Sects. \\ref{photo_pipe} and \\ref{spectro_pipe} we discuss how the photometric calibration was established for the photometer and spectrometer.  In Sect. \\ref{testing_models} we present the results of comparing SPIRE observations to solar system objects with known or well constrained model flux densities.  The beamsize and astrometric accuracy are discussed in Sect. \\ref{beam} and the wavelength accuracy in Sect. \\ref{wavelength_calibration}.  We summarise the present calibration status of SPIRE and draw some conclusions for future calibration activities in Sect. \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions} The most significant part of calibrating an astronomical instrument is the accurate conversion from measured signal to source flux density. In this paper we have given an overview of how this is currently carried out for the \\it{Herschel}\\rm-SPIRE instrument.  We have demonstrated an inter-comparison between the models used for astronomical sources and the \\it{Herschel}\\rm~telescope emission spectrum that gives a final agreement between the Neptune model used as the primary calibration standard and our measured spectrum of between 5-10\\%.  The calibrator source model is only one component of the final uncertainty and we have tested the conversion parameters for the photometer Volts to Jy using a Monte Carlo approach to determine the variance in the calibration curve as a function of voltage. This method involved creating multiple calibration curves of $1/\\Delta V_{PCAL}$ versus $V$ with addition random noise representative of the measurement errors and measuring the effect on the uncertainty in derived calibration coefficients. This translates into uncertainties that are $\\leq$ 0.1\\% for all detectors.  This is much smaller than the other two sources of error, which are the uncertainty in the model for Neptune (5-10\\%) and the uncertainty in peak flux density measurements for Neptune derived from the maps after processing. Taking these factors into account for data processed with the current processing software the overall photometric error can be assumed to be 15\\%.  For the spectrometer photometric calibration a comparison of the pipeline derived Neptune spectrum based on Vesta to the Neptune model shows a typical variation of 15-20\\% in the SSW band and 20-30\\% in the SLW band above 20 cm$^{-1}$ and up to 50\\% below 20 cm$^{-1}$.  It should be noted that variations in the temperature of the instrument itself between the measurement of the dark spectrum and the source can cause extra flux to be observed in the SLW band. The net result is an additive term in the spectrum below 25 cm$^{-1}$ of the order of a few Jy which can leave significant positive or negative extraneous flux density levels in the source spectrum especially for faint objects. With careful interactive processing it is possible to remove this signature and obtain accurate source flux densities (Ivison et al. \\cite{ivison10}, Bocklee-Morvan et al. \\cite{bockl10}). Further investigation on this aspect of the spectrometer calibration is ongoing and will be reported in future work.  The calibration of the data from the SPIRE instrument processed through the current data reduction pipeline software is already of a standard more than sufficient to allow significant scientific conclusions to be drawn from the observations.  Future work will concentrate on improving the photometry using other astronomical sources and source models, including cross calibration with other facilities such as the \\it{Planck}\\rm~satellite.   We will also work on improvements to the data processing tools, the astrometry of the maps through better attitude reconstruction and the wavelength calibration of the spectrometer through more sophisticated instrument performance models."
    },
    "1005/1005.2186_arXiv.txt": {
        "abstract": "CoRoT-7 b is the first confirmed rocky exoplanet, but, with an orbital semi-major axis of 0.0172 AU, its origins may be unlike any rocky planet in our solar system. In this study, we consider the roles of tidal evolution and evaporative mass loss in CoRoT-7 b's history, which together have modified the planet's mass and orbit. If CoRoT-7 b has always been a rocky body, evaporation may have driven off almost half its original mass, but the mass loss may depend sensitively on the extent of tidal decay of its orbit. As tides caused CoRoT-7 b's orbit to decay, they brought the planet closer to its host star, thereby enhancing the mass loss rate. Such a large mass loss also suggests the possibility that CoRoT-7 b began as a gas giant planet and had its original atmosphere completely evaporated. In this case, we find that CoRoT-7 b's original mass probably didn't exceed 200 Earth masses (about 2/3 of a Jupiter mass). Tides raised on the host star by the planet may have significantly reduced the orbital semi-major axis, perhaps causing the planet to migrate through mean-motion resonances with the other planet in the system, CoRoT-7 c. The coupling between tidal evolution and mass loss may be important not only for CoRoT-7 b but also for other close-in exoplanets, and future studies of mass loss and orbital evolution may provide insight into the origin and fate of close-in planets, both rocky and gaseous. ",
        "introduction": "\\indent The recently discovered exoplanet CoRoT-7 b has a mass of $M_p = 4.8 \\pm 0.8$ Earth masses ($M_{Earth}$) and a radius of only $R_p = 1.68 \\pm 0.09$ Earth radii ($R_{Earth}$). Thus its density likely lies between 0.71 and 1.4 times Earth's density, which suggests CoRoT-7 b is made of similar materials \\citep{2009A&A...506..287L, 2009A&A...506..303Q}. It appears to be the first example of a rocky exoplanet. However, with a semi-major axis $a$ of only $0.0172 \\pm 0.00029$ AU and $e \\approx 0$ \\citep{2009A&A...506..287L}, CoRoT-7 b has an orbit unlike the rocky planets in our solar system. CoRoT-7 b's surface may even be partially melted as a result of the strong stellar insolation \\citep{2009A&A...506..287L}. (Consequently, when we refer to CoRoT-7 b's ``rocky'' surface, we are referring to the composition and do not mean to imply that the surface is solid rock.) The announcement of an additional, nearby planet in the system, CoRoT-7 c, ($M_p = 8.4 \\pm 0.9$ $M_{Earth}$, $a = 0.046$ AU, and $e \\approx 0$) \\citep{2009A&A...506..303Q} also suggests that the system has an interesting dynamical history. \\\\ \\indent Missions like \\emph{Kepler} \\citep{2008IAUS..249...17B} and \\emph{CoRoT} \\citep{2006cosp...36.3749B} may find dozens of additional rocky exoplanets through transit observations. Although Kepler is expected to find several rocky planets in orbits like the Earth's \\citep{2008IAUS..249...17B}, the first and most easily detectable rocky planets will have orbits closer to their stars \\citep{2009PASP..121..952D}. Thus a detailed understanding of such planets will be critical to understanding the planets first discovered by these space-based missions. \\\\ \\indent The origins of these planets may be very different from the origins of rocky planets in our solar system. It may be difficult for rocky exoplanets to form as close to their stars as CoRoT-7 b and c because high temperatures in the protoplanetary gas disk may inhibit condensation and accretion of solid materials (\\emph{e.g.},  \\citealt{1994ApJ...421..615P}). At the very least, solid planets that form close-in will probably be depleted in volatile materials (\\emph{e.g.}, water [\\citealt{2008MNRAS.384..663R}]). Instead of forming in their current orbits, close-in rocky exoplanets like CoRoT-7 b may have coalesced in orbits farther away from their host stars, where disk temperatures were lower, and were brought to their current orbits. A variety of processes for bringing planets to such close-in orbits have been proposed \\citep{2004ApJ...614..955M, 2007Sci...318..210G, 2008MNRAS.384..663R}, and the orbital architecture of systems with close-in planets and the planets' physical properties provide clues to their histories \\citep{2008MNRAS.384..663R}. \\\\ \\indent One process that may have been important during CoRoT-7 b's history is evaporative mass loss. Evaporation is thought to influence gaseous exoplanets made primarily of H and He. For example, observations of HD 209458 b suggest the planet's atmosphere is evaporating at a rate of order 0.1 $M_{Earth}$/Gyr \\citep{2003Natur.422..143V,2004ApJ...604L..69V}, in line with theoretical expectations from studies of evaporative mass loss \\citep{2003ApJ...598L.121L, 2004A&A...418L...1L, 2004Icar..170..167Y, 2005ApJ...621.1049T, 2008SSRv..139..437Y}. Moreover, gaps in the distributions of orbital and physical properties of close-in planets have also been interpreted to indicate that many gaseous planets have had their atmospheres completely evaporated \\citep{2009MNRAS.396.1012D}. The complete removal of a gas giant's atmosphere would likely leave behind its rocky core, with a mass of perhaps several $M_{Earth}$ \\citep{2004A&A...419L..13B, 2008MNRAS.384..663R}. These considerations suggest there may be a large population of close-in planetary bodies that are remnant cores of evaporated gas giants \\citep{2005A&A...436L..47B, 2008PhDT........15H}. On the other hand, other studies argue that the observations of HD 209458 b do not imply significant loss of mass \\citep{2007ApJ...671L..61B, 2008Natur.451..970H}, and a theoretical study by \\citet{2009ApJ...693...23M} suggested that complete evaporation of a gas giant's atmosphere is unlikely. These competing hypotheses may now be testable. With the detection capabilities of the \\emph{Kepler} and \\emph{CoRoT} missions, rocky planets arising from a variety of histories may be detected. \\\\ \\indent Evaporation of mass may continue even after the planet loses its original H/He envelope. \\citet{2009arXiv0907.3067V} suggest the strong insolation received by CoRoT-7 b may be sufficient to have evaporated and removed several Earth masses of material, if the planet has always been solid (made of ice and/or rock). \\citet{2009ApJ...703L.113S} showed that thermal evaporation of rocky material from CoRoT-7 b's surface may produce a tenuous exosphere, similar to Mercury's. In this study, we consider evaporation both of a putative gaseous envelope and of solid material during CoRoT-7 b's history, including effects of its evolving orbit.\\\\ \\indent In addition to mass loss, tides have played an important and interrelated role in CoRoT-7 b's history. Tides have shaped the distribution of orbital elements for close-in exoplanets, reducing eccentricities \\citep{1996ApJ...470.1187R, 2006ApJ...638L..45F, 2008ApJ...678.1396J} and semi-major axes \\citep{2009MNRAS.395.2268B, 2009ApJ...698.1357J, 2009ApJ...692L...9L}. Tides may also help to bring rocky planets into close-in orbits \\citep{2008MNRAS.384..663R}. In fact, going forward in time, CoRoT-7 b may migrate into its host star in less than a few billion years, and it may have undergone past orbital migration \\citep{2009ApJ...698.1357J}. Also, CoRoT-7 b likely has undergone or currently undergoes strong tidal heating, either because its original orbit was eccentric or interactions with nearby CoRoT-7 c keep the eccentricity non-zero \\citep{2010ApJ...709L..95B}. Tidal evolution of its orbit may have been even faster in the past if CoRoT-7 b once had a massive gaseous envelope because the rate of evolution of a circular orbit scales with its mass. As the planet lost mass, though, the tidal evolution may have slowed as the planet reached its current orbit.\\\\ \\indent In this paper, we study the coupling of evaporative mass loss and tidal evolution on CoRoT-7 b. Using parameterized models for mass loss and orbital decay due to tides, we evolved CoRoT-7 b's mass and orbit backward (and forward) in time to determine the range of original masses and orbits consistent with the planet's current mass and orbit. Although the exact initial conditions depend sensitively on the chosen model parameters, we considered a wide range of parameters to map out all the possibilities. We find that, if CoRoT-7 b has always been a solid planet, it may have lost as much as half its original mass and had its semi-major axis reduced by as much as 60\\% by tidal migration alone. If CoRoT-7 b was originally a gas giant planet and only recently lost its gaseous envelope, its original mass did not exceed 200 $M_{Earth}$, and its original semi-major axis could have been almost twice as large as the current value. In both scenarios, CoRoT-7 b may have encountered and migrated through a 3:1 mean motion resonance with nearby CoRoT-7 c, as tides caused planet b's orbit to decay. Consideration of the resulting dynamical interactions may place constraints on the orbital history of the CoRoT-7 system and the original mass of CoRoT-7 b. We also discuss how the coupling between orbital decay and mass loss may play an important in the evolution and survival of gaseous close-in planets. ",
        "conclusions": "\\indent The first discovered rocky exoplanet, CoRoT-7 b, could actually be the remnant core of a gas giant. The planet's proximity to its host star and strong stellar insolation suggest that the planet may have lost a significant amount of mass over its lifetime. The proximity also suggests tides have played an important role in the planet's history, causing it to migrate inward. By coupling models of evaporative mass loss and tidal evolution, we find that the interplay between orbital decay and mass loss can lead to a wide variety of evolutionary pathways. Using a model that couples these two processes, we obtain constraints on the masses and orbits of gas giants that may have given rise to the CoRoT-7 b we see today. Our results suggest a complex orbital history for the planets in the CoRoT-7 system, possibly including passage through mean-motion resonances, even if CoRoT-7 b has always been a rocky planet.\\\\ \\indent Our calculations have general implications for the histories of both close-in rocky and gaseous planets. Previous studies have pointed out the signatures of evaporation and tidal evolution on the distribution of masses and semi-major axes for gaseous close-in exoplanets, but these studies have not considered the coupling of the two processes. Although our calculations allow CoRoT-7 b to be the remnant rocky core of a gas planet, confirming this hypothesis will require additional evidence. For example, improved constraints on the compositions of gas giant cores might allow us to distinguish remnant rocky cores on the basis of their physical properties. More detailed dynamical studies of CoRoT-7 and other similar systems will provide important constraints on their origins and histories. Whatever the details, mass loss and tidal evolution both probably played key roles in the histories and origins of many close-in exoplanets, rocky or gaseous.\\\\"
    },
    "1005/1005.3540_arXiv.txt": {
        "abstract": "We present mid-infrared spectra of six FeLoBAL QSOs at $1<z<1.8$, taken with the {\\it Spitzer} space telescope. The spectra span a range of shapes, from hot dust dominated AGN with silicate emission at 9.7$\\mu$m, to moderately obscured starbursts with strong Polycyclic Aromatic Hydrocarbon (PAH) emission. The spectrum of one object, SDSS 1214-0001, shows the most prominent PAHs yet seen in any QSO at any redshift, implying that the starburst dominates the mid-IR emission with an associated star formation rate of order 2700 M$_{\\odot}$ yr$^{-1}$. With the caveats that our sample is small and not robustly selected, we combine our mid-IR spectral diagnostics with previous observations to propose that FeLoBAL QSOs are at least largely comprised of systems in which (a) a merger driven starburst is ending, (b) a luminous AGN is in the last stages of burning through its surrounding dust, and (c) which we may be viewing over a restricted line of sight range. ",
        "introduction": "It is now well-established that there exist intimate links between the growth of supermassive black holes and the production of stars in galaxies. Indirect evidence for such links comes from, for example, the tight relationship between central black hole and stellar bulge masses (e.g. \\citealt{mag98,geb00,gad09}), and the coeval stellar populations in many massive elliptical galaxies \\citep{dun96,ell97,rak07}, suggesting their stars formed within less than a Gyr of each other. Direct evidence comes from the existence of galaxies that simultaneously harbor both high rates of star formation (of order tens to thousands M$_{\\odot}$yr$^{-1}$) and rapid accretion onto a central black hole \\citep{san96,gen98,far03,lon06}, all enveloped in large quantities of dust, making them extremely luminous in the infrared. The importance of these links for understanding the assembly history of galaxies over at least a substantial fraction of the age of the Universe is demonstrated by the strong evolution of the luminosity function of IR-luminous galaxies with redshift \\citep{sau90,lef05}, and the existence of a profusion of IR-luminous sources in the high-redshift Universe (e.g. \\citealt{bar98,eal99,cop06,aus09}).  An insightful way to study these links is to identify active galaxies at `key' points, where they are either rapidly building up stellar or central black hole mass, or shifting from one evolutionary phase to the next. One such point may be the \"youthful\" QSOs; if the starburst precedes or is coeval with the QSO\\footnote{The (optical) QSO phase is not necessarily the period in which the SMBH gains the bulk of its mass; indeed, there is evidence that the major periods of BH growth are coeval with or precede the starburst, see e.g. \\citet{mar05,mer10,trei10}}, then young QSOs should be those in which the starburst is drawing to a close, and a QSO is starting to emerge from its surrounding dust. Such objects would provide an excellent laboratory for testing feedback mechanisms between the starburst and AGN (e.g. \\citealt{sil98,cio07,lag08,moe09,cev09}).  Finding the youthful QSO population is a challenging and ongoing problem (e.g. \\citealt{san88,can01,lac06,cop08,geo09,lip09}). Significant attention has focused on the Broad Absorption Line (BAL) QSOs \\citep{lyn67,wey91,bro98,sch99,ara01,gre01,hall02,rei03,pri07,gib09}, whose properties\\footnote{BAL QSOs come in 3 subtypes. High Ionization BAL QSOs (HiBALs) show absorption in CIV $\\lambda$1549, NV $\\lambda$1240, SiIV $\\lambda$1394 and Ly$\\alpha$, and comprise about 10\\%\\ of all QSOs \\citep{tru06}. Low Ionization BAL QSOs (LoBALs) additionally show absorption in MgII $\\lambda$2799 and other lower ionization species, and comprise $\\sim1.3$\\% of all QSOs. Finally, FeLoBAL QSOs, in addition to showing all the absorption lines seen in LoBALs, also show weak iron absorption features \\citep{haz87,bec97}. They account for around $0.33\\%$ of all QSOs, though the exact fraction is unclear.} have been explained either as arising from being observed at a particular orientation, or because they are young objects still partially surrounded by dust. The properties of HiBALs are now thought to be primarily an orientation effect \\citep{surd87,mur95,sch99,gal07,doi09}, while the LoBALs remain controversial, with evidence favoring both orientation and evolution (e.g. \\citealt{voi93,ogl99,gal07,gho07,mon08,urr09,zha10}). For FeLoBALs there is also controversy, but the evidence favoring the evolution scenario is (arguably) stronger than for the LoBALs. Based on rest-frame UV and optical spectra, Hall et al (2002) suggest that FeLoBALs are young objects still enveloped in dust. Similar conclusions are reached by Gregg et al (2002), who also postulate that FeLoBALs may be associated with mergers. Further evidence comes from the discovery of iron absorption lines in the UV spectra of two low-redshift ULIRGs which harbor obscured AGN \\citep{far05}. Finally, mid-IR photometry observations (\\citealt{far07a}, hereafter F07) found that many FeLoBAL QSOs were IR-luminous, and may harbor high rates of star formation.  Excellent distinguishing evidence of the `youth vs. orientation' debate for FeLoBAL QSOs would be observations that show directly that high rates of star formation accompany the AGN. Given their reddened nature, a good place to look for such evidence is in the md-infrared, in which the spectral shapes of reddened AGN differ markedly from those of star-forming regions. The outstanding capabilities of the Infrared Spectrograph (IRS, \\citealt{hou04}) on-board the Spitzer space telescope \\citep{wer04} provided a huge step forward in available mid-IR spectroscopic capabilities, with the opportunity to shed light on the FeLoBAL phenomenon. In this paper, we use Spitzer-IRS observations of six FeLoBAL QSOs to examine the idea that they are young objects. We assume a spatially flat cosmology, with $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega=1$, and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "A general caveat to all of what follows is our small and heterogeneous sample. Our conclusions should thus be regarded as tentative.  The mid-IR spectra of our sample span a wide range of shapes. We see classical QSO spectra with hot silicate dust, together with classical starburst spectra with strong PAHs. The `starburst' spectra have more prominent PAHs than those seen in any other QSO so far observed, while the hot silicate dust sources do not have a close match in any purely mid/far-IR selected sample that we are aware of. It is possible that our heterogeneous selection leads to the heterogeneity of the spectra, but if this were true then we might expect the ISO selected object to have the strongest PAHs, which is not the case. Indeed, the two spectra at the extreme ends of the range of spectral shapes are both SDSS objects. The star formation rates of our sample span levels comparable to those in the most luminous starbursts found in any survey at any wavelength, to those seen in moderately IR-luminous AGN. The spectral diagnostics mark our sample as late-stage ULIRGs with a wide range of fractional AGN contributions.  These results, combined with their optical classification as reddened QSOs with a strong outflow, are in principle compatible with a continuum of `end-of-ULIRG' vs `peculiar QSO' contributions to the FeLoBAL QSO population. To frame the discussion, we describe three points on this continuum in detail:  \\noindent {\\bf 1:} FeLoBAL QSOs are a transition stage between starburst dominated ULIRGs and QSOs. As our sample span the entirety of this transition, and because FeLoBALs are observed to be rare\\footnote{A further reason behind the rarity of FeLoBALs is that they they are hard to detect; FeLoBAL QSOs are red, have few or no broad emission lines and no UV excess, and are thus hard to find in standard QSO searches (see e.g. \\citealt{app05}). The iron absorption features themselves are weak.}, we can set constraints on the timescale of this transition depending on how intrinsically common the Fe absorption features are. If Fe absorption is ubiquitous in such a transition, then the transition must be short in comparison to the ULIRG or QSO lifetime. Taking the ULIRG and QSO lifetimes to be $\\sim$10$^{8}$ years \\citep{tac08}, then the transition must take $\\lesssim$10$^{7}$ years (see also \\citealt{gre02}). If on the other hand the iron absorption features are intrinsically rare, then the transition can take longer.  \\noindent {\\bf 2:} FeLoBALs are comprised of comparable fractions of two galaxy populations, one that is transitioning from a ULIRG to a QSO, and another that is an unusual phase that QSOs go through, unconnected to mergers. Here the heterogeneity of our sample arises from observing two unconnected populations.  \\noindent {\\bf 3:} FeLoBAL QSOs are entirely an `unusual QSO' class, and have nothing to do with mergers or otherwise `transitioning' systems. Instead, the iron absorption features mark QSOs with both an outflow and atypically large amounts of iron in their ISM.  Formally, we cannot rule any of these scenarios out. The third though seems unlikely. It is hard to see how such a heterogeneous set of mid-IR spectra could arise from observing a peculiar class of QSO in which nothing more interesting than an outflow with large quantities of iron is going on, especially if we assume that the restrictive selection on rest-frame UV properties means we are viewing them over a restricted range in line of sight. Furthermore, most systems with star formation rates in the range seen in our sample are ULIRGs, which (at low redshift, at least) are nearly all mergers.  Distinguishing between the first and second scenarios is however more difficult. On one hand, the first scenario uses a single origin for which there is independent evidence from studying ULIRGs. The second scenario thus seems contrived, as it proposes two origins where one will do. On the other hand, our sample is small, and it is tempting to place undue emphasis on the properties of SDSS 1214-0001, as it is so peculiar. If this object were removed then the `unusual QSO' scenario would become more attractive. That said, the counter-argument also holds; remove (say) SDSS 1556+3517, and the `end-of-ULIRG' scenario becomes more attractive. We are mindful though that it is only the properties of SDSS 1214-0001 (and to an extent SDSS 1427+2709) that make us seriously consider the `end-of-ULIRG' scenario.  Recent work that may shed light on this is the study of NLS1's by \\citet{san10}. Their spectra show that at least some NLS1's harbor intense star formation and a weak mid-IR AGN continuum. NLS1's also have strong optical Fe emission lines, a large soft X-ray excess, and exhibit rapid, large amplitude X-ray variability. Furthermore, their black hole masses appear to be smaller than those in BLS1's of comparable luminosity, though there is controversy over whether NLS1's lie below \\citep{pet00,gru04} or on \\citep{bot05,kom07,dec08} the $M_{BH} - \\sigma_{*}$ relation \\citep{tre02}. On one level this hints at interesting links between NLS1s and FeLoBAL QSOs; for example, FeLoBAL QSOs may be analogues of NLS1s seen more pole-on than edge-on, and in which a strong outflow has formed. Exploring this idea in detail requires significantly larger samples, so we do not pursue it here. More generally however, it implies that relative orientation {\\it may} also play a role in the FeLoBAL QSO phenomenon.  We think it likely that the simplest solution that is consistent with all the observations, and the apparent rarity of FeLoBALs, is the correct one. We therefore propose that FeLoBAL QSOs as a class are rapidly evolving, youthful QSOs in which a starburst is coming to an end, and the AGN is in the last stages of burning through its surrounding dust. We also propose, with more reserve, that (1) we view FeLoBAL QSOs over a restricted line of sight range, where we see them more pole-on than edge on, and (2) the outflow has in some cases partially cleared the dust from around the AGN, leaving the IR emission from the starburst to dominate.  There are three ways to test this conclusion. First, FeLoBAL QSO hosts should be in the final stages of merging, so their host galaxies should show slight signs of morphological disturbance. Second, larger mid-IR spectroscopic surveys of FeLoBAL QSOs should show a similar range of spectral shapes to ours. Third, far-IR photometric surveys of FeLoBALs should find both a moderate enhancement of the fraction of far-IR bright sources compared to the classical QSO population, corresponding to those FeLoBAL QSOs that still harbor high star formation rates, and a correlation between far-IR luminosity and PAH equivalent width. Moreover, it would be useful to perform radiative transfer modelling of a BAL wind through a starburst to determine the conditions under which iron absorption is observed in such systems, and so provide insight into whether our conclusions are feasible.  We close with a caveat and a speculative comment. First, there is an important parameter that we cannot address; the way in which iron absorption features in FeLoBALs signpost starburst and AGN activity. For example, if the Fe absorption only occurs when a starburst is ending {\\it and} an AGN is blowing away its surrounding dust, but can occur throughout such a transition, then our first scenario is likely correct. As a contrived counter-example, if Fe absorption can only occur during the early phases of such a transition, but can occur randomly in reddened AGN in which there is no significant star formation, then in a small sample such as ours we would conclude, incorrectly, that all FeLoBAL QSOs were a ULIRG-to-QSO transition where the Fe absorption is present throughout. Resolving this issue is beyond the scope of this paper, so we note it as a caveat to our conclusions.  Second, \\citet{hall02} note that two of our sample, SDSS 1214-0001 and SDSS 2215-0045 have [FeIII]/[FeII] ratios greater than unity, implying the BALs arise in unusually dense, hot gas. The two ratios however are different; SDSS 2215-0045 has a [FeIII]/[FeII] ratio much greater than unity, while SDSS 1214-0001 has an [FeIII]/[FeII] ratio only slightly greater than unity. While conclusions drawn on only two objects are not trustworthy, this difference in ratio is consistent with our proposed evolutionary sequence. SDSS 1214-0001 has strong PAHs and silicates in absorption, while SDSS 2215-0045 has weak PAHs and silicates in emission, implying that SDSS 1214-0001 is at an earlier evolutionary phase than SDSS 2215-0045. Moreover, the enhanced [FeIII]/[FeII] ratio in SDSS 2215-0045 compared to SDSS 1214-0001 could be interpreted as the BAL outflow in SDSS 2215-0045 being more developed\\footnote{This is consistent with the `youth' hypothesis that \\citet{hall02} give for SDSS 2215-0045}. Further support for this idea would be an [FeIII]/[FeII] ratio greater than unity in SDSS 1427+2709, but the only published spectrum we are aware of is in \\citet{bec00}, in which the FeIII lines lie in a noisy region at the extreme blue end of the bandpass. They do not however appear to be significantly stronger than the FeII lines, so we simply note this as an interesting avenue for future work."
    },
    "1005/1005.3295_arXiv.txt": {
        "abstract": "{The Abell 222 and 223 clusters are located at an average redshift z$\\sim $0.21 and are separated by 0.26~deg. Signatures of mergers have been previously found in these clusters, both in X-rays and at optical wavelengths, thus motivating our study.  In X-rays, they are relatively bright, and Abell~223 shows a double structure. A filament has also been detected between the clusters both at optical and X-ray wavelengths.  } {We analyse the optical properties of these two clusters based on deep imaging in two bands, derive their galaxy luminosity functions (GLFs) and correlate these properties with X-ray characteristics derived from XMM-Newton data.} {The optical part of our study is based on archive images obtained with the CFHT Megaprime/Megacam camera, covering a total region of about 1~deg$^2$, or 12.3$\\times$12.3~Mpc$^2$ at a redshift of 0.21. The X-ray analysis is based on archive XMM-Newton images.} {The GLFs of Abell~222 in the $g'$ and $r'$ bands are well fit by a Schechter function; the GLF is steeper in $r'$ than in $g'$. For Abell~223, the GLFs in both bands require a second component at bright magnitudes, added to a Schechter function; they are similar in both bands. The Serna \\& Gerbal method allows to separate well the two clusters. No obvious filamentary structures are detected at very large scales around the clusters, but a third cluster at the same redshift, Abell~209, is located at a projected distance of 19.2~Mpc. X-ray temperature and metallicity maps reveal that the temperature and metallicity of the X-ray gas are quite homogeneous in Abell~222, while they are very perturbed in Abell~223.} {The Abell~222/Abell~223 system is complex. The two clusters that form this structure present very different dynamical states. Abell~222 is a smaller, less massive and almost isothermal cluster. On the other hand, Abell~223 is more massive and has most probably been crossed by a subcluster on its way to the northeast.  As a consequence, the temperature distribution is very inhomogeneous. Signs of recent interactions are also detected in the optical data where this cluster shows a ``perturbed'' GLF.  In summary, the multiwavelength analyses of Abell~222 and Abell~223 are used to investigate the connection between the ICM and the cluster galaxy properties in an interacting system.} ",
        "introduction": "Galaxy evolution is known to be influenced by environmental effects, which are particularly strong in galaxy clusters, all the more when they are undergoing merging events. The analysis of galaxy luminosity functions (GLF), and in particular of their faint-end slopes is a good way to trace the history of the faint galaxy population and the influence of merging events on this population. As summarized for example in Table~A.1 of Bou\\'e et al. (2008), this slope can strongly vary from one cluster to another, and from one region to another in a given cluster.  Superseding the dynamical analyses of Proust et al. (2000) and Girardi \\& Mezzetti (2001), Dietrich et al. (2002) have obtained a redshift catalogue for 183 galaxies in the region of the cluster pair Abell~222/223, out of which 153 belong to one of the two clusters. Their analysis was also based on photometry in the V and R bands.  They have estimated redshifts of 0.2126$\\pm$0.0008 and 0.2079$\\pm$0.0008, and velocity dispersions of 1014 and 1032~km~s$^{-1}$, for Abell~222 and Abell~223 respectively. As indicated by Dietrich et al. (2002) these velocity dispersions are somewhat higher than those expected from the X-ray luminosities. Abell~222 appears relatively relaxed, with a galaxy velocity histogram consistent with a gaussian. On the other hand, the adaptive kernel galaxy density map of Abell~223 shows two peaks, though surprisingly the Dressler-Schectman test (Dressler \\& Schectman 1988) does not find these two peaks and the DIP statistics cannot reject the null hypothesis of a unimodal distribution, implying that a Gaussian parent population cannot be excluded.  Dietrich et al. (2002) derived respective mass to light ratios of 202$\\pm$43 and 149$\\pm$33 h$_{70}$ $M_\\odot/L_\\odot$ and luminosity functions in the R band down to absolute magnitudes of $-19.5$, and found that a bridge of galaxies seems to exist between the two clusters. Note that such bridges of galaxies joining two clusters are still quite rare, another example at a comparable redshift being the filament detected between Abell~1763 and Abell~1770 by Fadda et al. (2008).  A weak lensing analysis of the Abell~222/223 system was performed by Dietrich et al. (2005), who found evidence for a possible dark matter filament connecting both clusters. They interpret the difference in redshift between Abell~222 and Abell~223 as implying a separation along the line of sight of 15$\\pm 3$ h$_{70}^{-1}$~Mpc. Their weak lensing surface mass density contours again show evidence for a double structure in Abell~223.  The existence of a 1.2~Mpc wide filament between both clusters was also found in X-rays by Werner et al. (2008).  We present here a muti-wavelength optical and X-ray analysis of the clusters Abell~222 and Abell~223. The optical part is based on imaging in the $g'$ and $r'$ bands obtained with the CFHT Megaprime/Megacam camera and retrieved from the Megapipe image stacking pipeline at CADC (Gwyn 2009). As seen in Table~\\ref{tab:obs}, the $r'$ band data are about 0.55~magnitude deeper than the $g'$ band data (by taking an average color $g'-r'=0.85$), and will be given priority in the morphological analysis; the $g'$ data will be used to select bright cluster members from color-magnitude diagrams, and also because the $g'$ band is more sensitive to star formation than the $r'$ band. These images are deeper than those of Dietrich by at least 0.5~mag (our raw galaxy counts start decreasing for $r'>24.5$ while Dietrich et al. say that their counts in the R band stop following a power law for R$>24$) , and therefore allow to probe the GLF to somewhat fainter magnitudes. Our X-ray analysis is based on archive XMM-Newton data.  We will consider hereafter that the two clusters have a common redshift of 0.21, for which Wight's cosmology calculator (Wright 2006) gives a distance of 1035~Mpc, a spatial scale of 3.427~kpc/arcsec and a distance modulus of 40.07 (assuming a flat $\\Lambda$CDM cosmology with H$_0=70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega _M=0.3$ and $\\Omega _\\Lambda =0.7$).  The paper is organised as follows. We describe our optical analysis in Section 2, and results concerning the galaxy luminosity function in Section 3. A search for substructures in both clusters based on the Serna-Gerbal method is described in Section~4.  The X-ray data analysis and results, including temperature and metallicity maps, are presented in Section~5.  An overall picture of this cluster pair is drawn in Section~6. ",
        "conclusions": "Understanding the interplay between galaxies and ICM is fundamental for further investigations of galaxy clusters. In this way, an important step is made with multiwavelength studies of these structures.  In order to understand better the connection between ICM and galaxies in merging systems we combined both X-ray and optical data to examine in detail the pair of clusters Abell~222/Abell~223. To do so, we combined spectroscopically computed 2D-maps with optical data to gain insight on the cluster history.  We verified that besides their proximity, these clusters show very different evolutionary stages that can be detected in optical and X-ray analyses. The Abell~223 data indicate that this cluster is dynamically perturbed when compared to its companion. Signs of recent merger(s) are detected at both wavelengths, meaning that the gas and the galaxies are still out of equilibrium. It is more likely that the signs of merger detected in Abell~223 are due to the interaction with the subcluster.  On the other hand, Abell~222 is a smaller cluster closer to a stable dynamical state than Abell~223.  As mentioned in Sect.~\\ref{Xrayana}, Abell~223 exhibits a remarkable blob of hotter gas at its southwest extremity. The gas of this region may have been compressed and heated by a recent shock that this cluster has been through. Another zone of higher temperature is elongated towards the north east-south western direction.  This is close to the direction that connects the two clusters and where a filament of low temperature was already detected between these clusters (Werner et al. 2008). This region does not appear in our 2D maps because there are not enough X-ray counts. However, the X-ray contours overlaid in the 2D maps show that the emission goes beyond the limits of the maps and towards the region where the filament is located.  The presence of only a small number of hot blobs and the low X-ray temperature of the filament tend to indicate that this system is in the first stages of a merger.  This hints at a scenario in which the two clusters are at the first stages of a collision.  These clusters are beginning to collide and from the optical counterpart analysis, the central region of Abell~223 already exhibits signs of this interaction presenting a higher number of star forming galaxies. So it seems probable that a small cluster has just crossed Abell~223 and is heading towards the north east.  On the other hand, Abell~222 does not present strong evidence of interaction other than the absence of a cool-core.  We thus summarize our main results below: \\begin{itemize} \\item{From the optical analysis we conclude that Abell~222 is smaller and less massive than Abell~223. For Abell~222, the GLF is quite well fit by a Schechter function in both the $r'$ and $g'$ bands. On the other hand, since Abell~223 is a double cluster, it presents a ``perturbed'' GLF. Abell~223 has more bright galaxies than Abell~222 and as a consequence a second component is required to fit the GLF at bright magnitudes. The GLFs of both clusters are comparable in the $r'$ band but in the $g'$ band the GLF is notably steeper in Abell~223 than in Abell~222. This is explained by a higher number of faint star forming galaxies present in Abell~223. Star formation is triggered by recent merger events. Such mergers disturbed the gas and its imprints can be observed in 2D X-ray maps.}  \\item{Both clusters appear to be elongated when traced by their cluster galaxy distributions. This is another evidence that these clusters are not in equilibrium.  However, the X-ray filament joining the two clusters does not seem to play a very important role in the evolution of cluster galaxies. This is probably explained by the low temperature of the gas, since ram-pressure is not very efficient below 2~keV.}  \\item{From the spectrally measured 2D X-ray maps we conclude that Abell~222 is fairly isothermal when compared to Abell~223. However, neither of these clusters presents a cool-core as a sign of relaxation. The overall temperature map of Abell~223 is an imprint of the recent dynamical interactions this cluster has experienced.  The broad variation in temperature puts in evidence a large number of sub-structures.  From the metallicity map, we also observe an enhancement in the central region of Abell~223. This can be explained by the fact that this cluster has more very bright galaxies than Abell~222. Assuming ram-pressure stripping with tidal disruption of the bright galaxies near the cluster center as a mechanism for metal injection, one would expect this cluster to have a higher central metallicity than Abell~222. The combined analysis of X-ray and optical data argues for a scenario in which the two clusters are beginning to collide. The former cluster presents signs of interaction in both wavelength ranges, and has probably been crossed by a small cluster heading north east, while its companion does not show strong evidence for interaction, except for the absence of a cool-core.}  \\end{itemize}"
    },
    "1005/1005.4683_arXiv.txt": {
        "abstract": "The next generation mass probes will obtain information on non--linear power spectra $P(k,z)$ and their evolution, allowing us to investigate the nature of Dark Energy. To exploit such data we need high precision simulations, extending at least up to scales of $k \\simeq 10\\Mpch$, where the effects of baryons can no longer be neglected. In this paper, we present a series of large scale hydrodynamical simulations for \\LCDM and dynamical Dark Energy (dDE) models, in which the equation of state parameter is $z$-dependent. The simulations include gas cooling, star formation and Supernovae feedback.  They closely approximate the observed star formation rate and the observationally derived star/Dark Matter mass ratio in collapsed systems. Baryon dynamics cause spectral shifts exceeding $1\\, \\%$ at $k > 2-3 h$Mpc$^{-1}$ compared to pure n-body simulations in the \\LCDM simulations.  This agrees with previous studies, although we find a smaller effect ($\\sim 50\\%$) on the power spectrum amplitude at higher $k$'s. dDE exhibits similar behaviour, even though the dDE simulations produce$\\sim 20\\, \\%$ less stars than the analogous \\LCDM cosmologies. Finally, we show that the technique introduced in Casarini \\etal to obtain spectra for any $w(z)$ cosmology from constant--$w$ models at any redshift still holds when gas physics is taken into account. While this relieves the need to explore the entire functional space of dark energy state equations, we illustrate a severe risk that future data analysis could lead to misinterpretation of the DE state equation. ",
        "introduction": "There can be little doubt that Dark Energy (DE) is a necessary ingredient for any cosmological model approaching data (see, {\\it e.g.}, the data fits in WMAP5 release, Komatsu \\etal 2009). In turn, the quest for the nature of DE is perhaps the main challenge in contemporary physics in spite of the fact that no available data conflict with \\LCDM models, in which the DE state parameter is $w \\equiv -1$, as is true for a false vacuum. A false vacuum interpretation of DE, however, leads to an unnatural value for its density (fine tuning problem) and to a number of coincidences characterizing our epoch and the primeval fluctuation amplitude (coincidence problem).  Many efforts have therefore been devoted to finding viable alternatives.  Among them, the possibility that DE is a scalar field, $\\phi,$ self-interacting through a potential $V(\\phi)$ (Wetterich 1989, Ratra \\& Peebles 1989), has been widely explored. In this case, both the ratio between kinetic and potential energies, $\\dot \\phi^2/2V(\\phi)$ and the DE state parameter $w(z)$, depend on redshift, $z$. In this paradigm, observational data on $w(z)$ would yield $V(\\phi)$.  We will refer to this class of models, as dynamical DE (dDE).  Future large tomographic shear surveys, such as the Euclid project (see, {\\it e.g.}, Refr\\'egier \\etal 2010), will measure matter density fluctuations and their evolution from the linear to the non-linear regime with unprecedented precision, approaching $\\approx 1\\%$ (Huterer \\& Takada 2005). To translate these data into information on DE, we need to predict power spectra and the evolution of power spectra with similar precision for a wide set of $w(z)$ laws.  The resulting constraints on $w(z)$ will be a substantial step toward understanding DE physics (see, {\\it e.g.}, Hu \\& Tegmark 1999, Dodelson \\& Zhang 2005, Manera \\& Mota 2005, Mota 2008, La Vacca \\etal 2008, 2009, Kristiansen \\etal 2010.).  The evolution of the baryonic component of the Universe is highly uncertain and can affect predictions for lensing observables, manifesting as modified structure growth for a fixed cosmic distance scale (Zentner et al 2008, Hearin \\& Zentner 2009).  Collisionless N--body simulations have been extensively used to obtain non-linear power spectra (e.g. Smith \\etal 2003).  These simulations have ignored the effect of baryons, which comprise $\\sim 15\\, \\%$ of the matter content of the observed Universe, (Komatsu \\etal 2009).  Although baryons make up a minor fraction of the Universal matter content and is distributed exactly the same as Cold Dark Matter (CDM) at the onset of the non-linear growth, the final distribution of baryons in haloes is significantly different than CDM, because of its separate later dynamical evolution.  Jing \\etal (2006) have shown that, in the spectra of \\LCDM models, the difference between N--body simulations and simulations including baryon physics exceeds 1\\% at $k > 2$--3$\\, h{\\rm Mpc}^{-1}, $ and increases to $\\sim 10$--20$\\, \\%$ on smaller scales when $k$ approaches 10$\\, h{\\rm Mpc}^{-1} $. Rudd \\etal (2008) and Guillet \\etal (2010) recently confirmed the need for simulations including baryonic physics, even though their results differed from Jing \\etal (2006).  The discrepant results originate from the different prescriptions used for star formation, SN feedback, AGN effects, and possibly, the numerical approach.  To compare different $w(z)$ laws, these uncertainties need to be under control. Using the same prescriptions for all models is essential to compare their spectra, as we do in this work. Their being tuned to observations is also a need, when aiming to fit future data. The prescription used here are tuned on available observations, such tuning being however carried on within the frame of $\\Lambda$CDM cosmologies. This is to be borne in mind, in view of model dependent effects that we shall detect, e.g. for what concerns star production in haloes.  Dynamical Dark Energy (dDE) simulations, with a variable state parameter $w(a)$ deduced from scalar field potentials admitting a tracker solution, have been performed since 2003 (e.g. Klypin \\etal 2003, Linder \\& Jenkins 2003, Dolag \\etal 2004, Macci\\`o \\etal 2004, Solevi \\etal 2005).  These have been compared with constant-$w$ simulation outputs. Observables considered in these papers, only marginally included spectra. Recently Francis \\etal (2007) have shown how spectral predictions for constant-$w$ models at $z=0$ can also be used to fit spectra of cosmologies with a state parameter given by a first degree polynomial: \\begin{equation} w(z) = w_0 + w'(1-a)~. \\label{wz} \\end{equation} More precisely, the spectrum $P(k,0)$, in an {\\it assigned} ($\\cal A$) model with state parameter $w(z)$, can be approached by a suitable {\\it auxiliary} constant-$w$ model ($\\cal W$) if: (i) in $\\cal A$ and $\\cal W$, $\\Omega_{b,m,tot}$, $h$ and $\\sigma_8$ are equal, (ii) the constant DE state parameter $w$ of $\\cal W$ is tuned so that $\\cal W$ and $\\cal A$ have the same comoving distance from the Last Scattering Band (LSB) and $z=0$.  In that case, spectral discrepancies stay $< 1\\, \\%$, up to $k \\sim 2$--$3\\, h\\, $Mpc$^{-1}$ (symbols here above have their usual meaning).  Francis \\etal (2007) also tried extending their results to higher redshifts, but discrepancies between $\\cal A$ and $\\cal W$ increased to several percent; not enough to exploit forthcoming data. The required precision was, however, recovered through a new technique introduced by Casarini, Macci\\`o \\& Bonometto (2009, Paper I hereafter) and tested with N--body simulations.  The final aim of this work is to extend the techniques use in Paper I to scales that only hydrodynamical simulation can test. This is a serious challenge for the description of physical processes included in hydrodynamical algorithms.  We will devote a part of our work to understanding these challenges.  The focus of this paper is twofold: (i) Study the importance of the effects of baryons on the matter power spectrum on the small scales that will probed by future surveys, bearing in mind that the physical assumptions will have to be carefully calibrated to approach the unprecedented precision requirements. (ii) Extend the commonalities found in $N$-body simulations in Paper I among spectra of different models down to scales where baryon physics cannot be disregarded. Accordingly, we shall use the same baryon physics in a series of simulations to compare the spectra of different dDE models.  The paper is therefore organized as follows: In \\S 2, we describe the approach used in Paper I. In \\S 3, we motivate the choice of the models considered. \\S 4 is devoted to describing our simulations and the techniques used to analyse them. In \\S 5, we present our hydrodynamical simulations. \\S 6 is subdivided in subsections that concern various aspects of the power spectra analysis. In \\S 6 we test: (i) the relation between N-body and hydrodynamical spectra (ii) the effect of box size and numerical resolution (iii) the level at which model parameters can be discriminated; and, finally, (iv) the success of the extension of Paper I approach to the higher $k$ range. A general discussion of the results closes our paper.   \\begin{figure} \\psfig{figure=wmap5.w.eps,width=0.52\\textwidth} \\caption{WMAP7 (solid curves) \\& WMAP5 (dashed curves) likelihood ellipses on the $w_o$--$w'$ plane and the position of our models on that plane. The significant shift from WMAP5 to WMAP7 is not due to new CMB data, but mostly from the improved HST determination of $H_0$. Dotted lines yield the best likelihood model. The selected model lies near the 1--$\\sigma$ contour for WMAP7, and within 2--$\\sigma$'s, for WMAP5 as well. Notice how much closer the $\\cal W$ models (empty polygons) are than the assumed cosmology (filled triangle). This highlights a serious observational risk (see \\S \\ref{sec:dis}). } \\label{ellypse} \\end{figure} ",
        "conclusions": "\\label{sec:dis}   In this work, we employed large scale hydrodynamical simulations, including gas cooling, star formation and Supernovae feedback.  Our simulations perform quite well when compared to observational results. They produce a sizable stellar component in collapsed dark matter haloes and star formation histories comparable with observations. We are aware that these results strongly depend on the adopted (low) resolution, nevertheless the similarity to observations makes us confident that the effect of baryons on the matter power spectrum in our simulation is sufficiently realistic. That said, further work is needed to better address the impact of the physics of star formation on the final results and, most important, to provide better tests of sample variance. In particular, this work makes clear that no really quantitative predictions, unaffected by sample variance, can be drawn from simulations, in $L\\approx 50 \\hMpc$ boxes. Matching observed accuracies of $\\cal O$$(1\\, \\%)$ will require boxes a few hundred Mpcs on a side. In principle a box with $L = 256\\, h^{-1}$Mpc, as the one we mostly used here, should be adequate to delve into haloes, providing spectra up to $k \\simeq 10\\, h\\, $Mpc$^{-1}$, with the required accuracy. However, the spectra obtained from our large box differed in a non negligible way from the ones obtained from a smaller box ($L=64 \\hMpc$). Although being fully compatible with the sample variance of the smaller box, this suggests further investigation will be required.  These residual uncertainties do not prevent us, though, from testing the efficiency of a method introduced in Paper I, that allows us to work out power spectra prediction for any DE model based on models with a constant equation of state ($w$). We found up to $z = 2$, that the {\\it auxiliary} $\\cal W$ models (with constant $w$) yield spectra consistent with our {\\it assigned} $\\cal A$ model (with a polynomial $w(z)$) within a few permils. The test was performed for a single $\\cal A$ dDE model. We plan to extend such a test to more dDE models in forthcoming work. There is however no reason to expect that the point selected on the $w_o$--$w'$ plane owns peculiar properties.  The residual discrepancies between $\\cal A$ and $\\cal W$ models exhibit a specific trend, already visible in Paper I: discrepancies (of a few permils) are greater for smaller $z$ values. Also DMG results follow the same rule, which is even even more pronounced for baryon spectra. Gas spectra discrepancies often exceed the permils but are systematically compensated by other components. The success of our technique was not unexpected, but discrepancies at such a tiny level, also for $k > 2$--$3\\, h\\, $Mpc$^{-1}$, exceeded our expectations.  The requirement that two models have the same $\\omega_{c,b}$ and $\\sigma_8 (z)$ introduces  basic similarity: they will have similar baryon acoustic oscillations and typical fluctuation amplitudes. Then, the requirement that $w$ yield the same conformal time $\\tau$ elapsed since recombination for assigned and auxiliary models apparently sets an equal timing for the gradual onset of the non-linear evolution on analogous structures. However, while $\\tau$ is a fair time coordinate for structures slightly detached from the cosmic background, it no longer plays such a role in cosmic islands, which have abandoned the cosmic expansion and evolve on their own, forgetting the continuing increase of the scale factor $a~.$ Within these structures, space-time is substantially Minkowskian, and time evolution is set by the ordinary time $t~.$  These are, however, two extreme situations. The scales which ``recently'' abandoned the overall expansion and entered a non-linear regime are in an intermediate position. Our inspection mostly concerns such scales. Accordingly, the success of an equal-$\\tau$ requirement is anything but granted.  Here we use hydrodynamical simulations to study a larger $k$ domain (compared to Paper I), where the effects of baryons cannot be neglected anymore. By introducing non-gravitational effects we expect dynamics to be regulated by the ordinary time $t$, thus producing larger discrepancies at scales $k > 2$--$3\\, h\\, $Mpc$^{-1}$. Figure \\ref{wr0} rejects this hypothesis by showing that the matter distribution is still ruled by the primeval clock. In particular, for $k > 2$--$3\\, h\\, $Mpc$^{-1}$, the overall DMG spectral discrepancies, at $z=0$, are perhaps even smaller than DMO. Curiously enough, the opposite seems true at some greater $z$'s, suggesting that most discrepancies found are pure noise, rather than systematics. Clearly for larger $k$'s, when we delve into regions that virialized long ago, $\\cal A$ and $\\cal W$ models unavoidably differ. However, while we restrict ourselves to the scales that forthcoming weak lensing surveys will study, the Paper I requirement still works unexpectedly well.   One ought to beware, however, that the success of our approach does not imply that constant-$w$ and variable-$w$ models are indiscernible.  What makes them observationally different is the measured dependence of $w$ on the redshift $z$.  An essential issue, however, is that {\\it the observed constant state parameter $w(z)$ will NOT be the physical one of the DE component}. Although we expect data to fit, at any $z$, the same density and Hubble parameter values, the observed $w(z)$ will be the state parameter that a constant-$w$ model must have, at such $z$, to grant the same distance between there and the LSB. The relation between the observed $w$ and the intrinsic one, for the DE component, should be put under control; otherwise there is a severe risk that future data analysis lead to DE state equation misinterpretation.  A final point we wish to make is that {\\sc halofit} expressions, so useful for predicting spectra until now, have been tested and found compatible with their claimed accuracy. Unfortunately, accuracy at ($\\sim 6\\, \\%$) no longer matches our needs. Furthermore, among models that approach the present data, {\\sc halofit} expressions only need to handle \\LCDM.  If one aims to improve {\\sc halofit} expressions, there is no longer any need to cover the wide array of models that once were relevant, but now are far from observation.  Amelioration, however, is needed in two separate directions: (i) The spectral range covered by predictions must reach $k \\simeq 10\\, h\\, $Mpc$^{-1}$, which requires that gas be considered; Rudd \\etal (2008) had actually attempted an extension in this direction, but limited themselves to \\LCDM cosmologies. (ii) DE state parameters different from constant $w \\equiv -1$ should be included.  Moreover, they should aim for a precision of $\\cal O$$(1\\, \\%)$, which presents a severe challenge, even for a much more restricted model range. Nevertheless when {\\sc halofit} expressions of the required precision will be available, even only for constant-$w$ models, thanks to the results of this work it will be straightforward to extend them to any  $w(z)$ law."
    },
    "1005/1005.1773_arXiv.txt": {
        "abstract": "{We use Herschel PACS and SPIRE observations of the edge-on spiral galaxy UGC\\,4754, taken as part of the H-ATLAS SDP observations, to investigate the dust energy balance in this galaxy. We build detailed SKIRT radiative models based on SDSS and UKIDSS maps and use these models to predict the far-infrared emission. We find that our radiative transfer model underestimates the observed FIR emission by a factor of two to three. Similar discrepancies have been found for other edge-on spiral galaxies based on IRAS, ISO, and SCUBA data. Thanks to the good sampling of the SED at FIR wavelengths, we can rule out an underestimation of the FIR emissivity as the cause for this discrepancy. Instead we support highly obscured star formation that contributes little to the optical extinction as a more probable explanation.} ",
        "introduction": "Edge-on spiral galaxies are an important class of galaxies in which the distribution and properties of interstellar dust grains can be studied in great detail. The dust in these systems shows prominently as dust lanes in optical images; for several edge-on spiral galaxies, the dust distribution has been modelled by fitting realistic radiative transfer models to such optical images \\citep{1987ApJ...317..637K, 1997A&A...325..135X, 1998A&A...331..894X, 1999A&A...344..868X, 2004A&A...425..109A, 2007A&A...471..765B}. The conclusion of these works is that, in general, the dust disc is thinner (vertically) but radially more extended than the stellar disc and that the central optical depth perpendicular to the disc is less than one in optical wavebands, making the disc almost transparent when seen face-on.  A complementary way of studying the dust component in galaxies is through its thermal emission at far-infrared (FIR) and submm wavelengths. A self-consistent treatment of extinction and thermal emission, i.e.\\ a study of the dust energy balance, gives the strongest constraints on the dust content of spiral galaxies. When we quantitatively compare the results from extinction studies and FIR/submm emission studies, a discrepancy is found: the relatively optically thin dust discs emerging from the optical modelling absorb about 10\\% of the stellar radiation, whereas FIR studies of normal spiral galaxies indicate they typically reprocess about 30\\% of the UV/optical radiation \\citep{2002MNRAS.335L..41P}. This problem is clearly illustrated when applied to individual edge-on spiral galaxies: the predicted FIR fluxes of self-consistent radiative transfer models that successfully explain the optical extinction generally underestimate the observed FIR fluxes by a factor of about three \\citep{2000A&A...362..138P, 2001A&A...372..775M, 2004A&A...425..109A, 2005A&A...437..447D}. Several scenarios (see Sect.~{\\ref{Results.sec}}) have been proposed to explain this discrepancy, but a major problem distinguishing between these is that the number of edge-on galaxies for which such detailed studies have been done so far is limited, owing to the poor sensitivity, spatial resolution, and limited wavelength coverage of the available FIR instruments. Herschel \\citep{Herschel} offers the possibility of studying the dust energy-balance in spiral galaxies in more detail than before. The combination of the sensitivity and the wavelength coverage of the PACS \\citep{PACS} and SPIRE \\citep{SPIRE} instruments, together covering the 70 to 500~$\\mu$m region where the emission from cold dust dominates, enables us to make reliable estimates of the total thermal emission of the interstellar dust.  \\begin{figure*} \\centering \\includegraphics[width=\\textwidth]{14644fg1.pdf} \\caption{The multi-wavelength view of UGC\\,4754. The top row shows the Herschel images, the bottom row a selection of optical and NIR images from SDSS and UKIDSS. The field-of-view of each image is $100\\arcsec \\times 100\\arcsec$.} \\label{Observations.pdf} \\end{figure*}    A modest number of large, nearby edge-on spiral galaxies are targetted by Herschel key programs, including the Very Nearby Galaxies program and the Herschel Reference Survey \\citep{2010PASP..122..261B}. Much promise, however, is offered by the Herschel Astrophysical TeraHertz Large Area Survey \\citep[H-ATLAS,][]{2009arXiv0910.4279E}, which will image a total area of 550~deg$^2$ with PACS and SPIRE at 100, 160, 250, 350, and 500~$\\mu$m. Based on the Revised Flat Galaxy Catalogue \\citep{1999BSAO...47....5K} and the 2MASS-selected Flat Galaxy Catalogue \\citep{2004BSAO...57....5M}, we expect between 100 and 200 resolved edge-on spiral galaxies in the H-ATLAS survey area.  In this Letter, we investigate the dust energy balance in the nearby spiral galaxy UGC\\,4754 as a first case. This Sbc galaxy is the nearest \\citep[$v_r=5662\\pm45$~km/s;][]{2009MNRAS.399..683J} and largest edge-on spiral galaxy in the 16~deg$^2$ field observed during the H-ATLAS science demonstration phase (SDP). We present detailed radiative transfer simulations of this galaxy that fit optical and near-infrared imaging data, compute synthetic FIR fluxes and images, and compare the model predictions to the Herschel observations at 100, 160, 250, 350, and 500~$\\mu$m. In Sect.~{\\ref{Observations.sec}} we present the observations and data reduction, in Sect.~{\\ref{Analysis.sec}} we present our radiative transfer modelling and compare the model results to the observations, and in Sect.~{\\ref{Results.sec}} we discuss these results and present out conclusions. ",
        "conclusions": "\\label{Results.sec}  An apparently straightforward explanation to the dust energy-balance problem might seem that the standard FIR emissivity, derived from dust models finetuned to the diffuse dust emission in the Milky Way, is actually underestimated by a factor two to three. This idea was advocated by \\citet{2000A&A...356..795A, 2004A&A...425..109A} and \\citet{2005A&A...437..447D} to explain the excess at SCUBA wavelengths of several edge-on spiral galaxies. Support for this explanation is the wide range of empirical values for the FIR emissivity that have appeared in the literature \\citep[see e.g.][]{1997MNRAS.289..766H, 2004A&A...425..109A}. Thanks to Herschel's wavelength coverage between 100 and 500~$\\mu$m, we can rule out this possibility, because we can accurately compare the absorbed stellar luminosity in our radiative transfer model ($L_{\\text{abs}} = 0.12~L_{\\text{bol}}$) to the observed luminosity $L_{\\text{d}}$ emitted by the dust. We calculated the observed dust luminosity $L_{\\text{FIR}}$ in the FIR (between 80~$\\mu$m and 1~mm) by integrating the SED of the best-fitting modified blackbody model to the Herschel data. This model, indicated as the green line in the bottom panel of Fig.~{\\ref{SED.pdf}}, corresponds to a modified blackbody model with $M_{\\text{d}} = 1.86\\times10^7~M_\\odot$ and $T=19.6$~K and yields $L_{\\text{FIR}} = 5.0\\times10^9~L_\\odot = 0.27~L_{\\text{bol}}$. Note that $L_{\\text{FIR}}$ is a lower limit to $L_{\\text{d}}$, since the mid-infrared luminosity (dominated by non-equilibrium emission by very small grains and PAHS) has not yet been taken into account. These numbers clearly indicate that there is a real discrepancy in the energy balance. This discrepancy cannot be lifted by changing the FIR emissivity at some wavelength; indeed, the energy balance equation impies that the ratio $L_{\\text{abs}}/L_{\\text{d}}$ is independent of the normalization of the emissivity.  The most likely explanation for this energy balance problem is that a sizable fraction of the FIR/submm emission arises from additional dust that has a negligible extinction on the bulk of the starlight. One option is the presence of a second, thinner dust disc, originally proposed by \\citet{2000A&A...362..138P}. Such a configuration has been successful in explaining the energy balance of spiral galaxies \\citep[][Popescu et al.\\ 2010, in prep.]{2001A&A...372..775M} and the observed attenuation-inclination relation \\citep{2007MNRAS.379.1022D}, but its validity has been questioned based on deep {\\em{K}} band images of NGC\\,891 \\citep{2005A&A...437..447D}. An alternative hidden dust source are young stars deeply embedded in dusty molecular clouds. The compact dust clumps can boost the FIR/submm emission of the dust, while keeping the extinction relatively unaltered \\citep[e.g.][]{1998ApJ...509..103S, 2000A&A...359...65B, 2005AIPC..761..155P, 2008A&A...490..461B, 2008ApJ...672..817M}.  A first indication that embedded star-forming clouds might be the solution to the case of UGC\\,4574 is that the discrepancy between our radiative transfer model and the observed FIR SED is stronger at shorter than at longer wavelengths. This implies that warmer dust (such as in star-forming regions) is necessary to bring the model in balance with the data. Further evidence is the forthcoming study by Smith et al.\\ (2010b, in prep.) who analyse the SED of 1260 galaxies from the H-ATLAS SDP field using the technique of \\citet{2008MNRAS.388.1595D}. This technique consists of fitting the observed SED with a combination of attenuated stellar emission, dust emission from star forming-regions, and dust emission from the general ISM. For UGC\\,4754, Smith et al.\\ (2010b) find that more than 50\\% of the bolometric luminosity of the galaxy is reprocessed by dust. In particular, their modelling predicts a significantly bluer intrinsic UV continuum. The very efficient absorption of this radiation in star-forming regions can power the FIR luminosity, while almost remaining unnoticed in large-scale optical extinction maps. Combining the results from that study with ours, we conclude that embedded star formation is the most likely way to lift the energy balance in UGC\\,4754.  The present case study of UGC\\,4754 is a demonstration of what we can learn from detailed dust energy balance studies of edge-on spiral galaxies. In future contributions we will apply similar modelling to larger sets of edge-on galaxies and investigate possible systematic links with other galaxy parameters; for example, one would expect a correlation between the strength of the dust energy balance problem and the star formation rate if embedded star formation is the main driver \\citep[e.g.][]{2008MNRAS.388.1595D, 2010MNRAS.403.1894D}. The Herschel observatory and, particularly, the H-ATLAS project offer the possibility to perform such studies in a systematic way.          \\begin{acknowledgement}  This work used data from the UKIDSS DR5 and SDSS DR7. The UKIDSS project is defined in \\citet{2007MNRAS.379.1599L} and uses the UKIRT Wide Field Camera \\citep[WFCAM;][]{2007A&A...467..777C}. Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. \\end{acknowledgement}"
    },
    "1005/1005.5016_arXiv.txt": {
        "abstract": "{We present {\\it Herschel} far-infrared (FIR) observations of two sub-mm bright quasars at high redshift: SDSS J1148+5251 ($z=6.42$) and BR 1202$-$0725 ($z = 4.69$) obtained with the PACS instrument. Both objects are detected in the PACS photometric bands. The {\\it Herschel} measurements provide additional data points that constrain the FIR spectral energy distributions (SEDs) of both sources, and they emphasise a broad range of dust temperatures  in these objects. For $\\lambda_{\\rm rest} \\lesssim 20$\\,$\\mu$m, the two SEDs are  very similar to the average SEDs of quasars at low redshift. In the FIR, however, both quasars  show excess emission compared to low-$z$ QSO templates, most likely from cold dust powered by vigorous star formation in the QSO host galaxies. For SDSS  J1148+5251 we detect another object at 160\\,$\\mu$m with a distance of $\\sim$\\,10\\arcsec~from the QSO.  Although no physical connection between the quasar and this object can be shown with the available data, it could potentially confuse low-resolution measurements, thus resulting  in an overestimate of the FIR luminosity of the $z = 6.42$ quasar. } ",
        "introduction": "The detection of large quantities of cold dust in high redshift ($z>5$) quasars \\citep[e.g.][]{ber03,bee06,wan08} implies a substantial enrichment of the interstellar medium already during the first billion years after the Big Bang. If the far-infrared (FIR) emission of these objects is powered by star-formation, then the luminosity and the temperature of the dust implies star formation rates of up to a few thousand solar masses per year, possibly indicating the rapid formation of early galactic bulges.  Recent observations at mid-infrared (MIR) wavelengths with {\\it Spitzer} demonstrated that the dust also shows an energetically important hot (T\\,$\\sim$\\,1000\\,K) component in many high-redshift QSOs \\citep{hin06,jia06}, with only few exceptions \\citep{jia10}. The similarity of the spectral energy distributions (SEDs) to those of lower redshift AGN lead to the conclusion that the general structures characterising local AGN are already in place at $z\\sim6$.  However, for a majority of the sources the (sub)mm and MIR observations have only measured the tail of the dust emission spectrum. In order to further constrain the properties of the dust in these objects, it is essential to sample the SED as completely as possible. These measurements can then be used to derive critical parameters such as the total infrared luminosity and the range of dust temperatures in the objects of interest. In addition, the MIR to FIR luminosity ratio may indicate the relative importance of warm and hot dust (which is predominantly heated by the AGN) compared to colder dust, which is preferentially heated by star formation.  The {\\it Herschel} key project \\citep{pil10} \"The Dusty Young Universe\" (PI K.~Meisenheimer) aims to measure the FIR SEDs for all quasars with $z>5$ that were known at the time of submission of the proposal (early 2007) with the PACS \\citep{pog10} and SPIRE \\citep{gri10} instruments. Together with existing NIR and MIR photometry, measurements in five {\\it Herschel} bands will yield complete infrared SED coverage of more than two decades in wavelength. Thus, the properties of the dust and the possible interplay between black-hole growth and galaxy bulge formation (traced indirectly through the FIR emission) can be explored early in the evolution of the universe.    \\begin{figure*}[t!] \\centering \\includegraphics[width=3.5cm]{14718f01a.ps} \\includegraphics[width=3.5cm]{14718f01b.ps} \\includegraphics[width=3.5cm]{14718f01c.ps} \\includegraphics[width=3.5cm]{14718f01d.ps}\\\\ \\includegraphics[width=3.5cm]{14718f01e.ps} \\includegraphics[width=3.5cm]{14718f01f.ps} \\includegraphics[width=3.5cm]{14718f01g.ps} \\caption{Infrared and (sub)mm images of J1148+5251. All images are 60\\arcsec~wide and north is up with east to the left. Another nearby source northwest of the QSO (see text) is clearly visible at longer wavelengths.} \\label{pic_j1148} \\end{figure*} ",
        "conclusions": "In Fig.~\\ref{sed_j1148} we show the spectral flux and energy distributions of J1148+5251 and 1202$-$0725.  The SEDs of both high-redshift QSOs appear very similar in shape, and the strong emission in all infrared  bands shows dust at a wide range of temperatures. This is particularly  emphasised by the {\\it Herschel} photometry, which fills the gap between the previously available MIR and sub-mm photometry.  The shape of the SED also implies that large amounts of dust are present in the host galaxies already at high redshift and that the dust may be distributed on a wide range of scales: NIR emission from hot dust close to the nucleus \\citep[e.g.][]{hin06,jia06,jia10}, warm MIR dust (T\\,$\\sim$\\,few hundred K) on intermediate scales (or partly shielded), and colder dust in the FIR possibly distributed throughout the host galaxy. Using the new {\\it Herschel} photometry to determine total infrared luminosities for these two objects (integrating the SEDs between 1 and 200\\,$\\mu$m, rest frame) yields  $L_{\\rm FIR} = 8.1 \\times 10^{13}\\,L_{\\odot}$ for J1148+5251 and a total $L_{\\rm FIR} = 3.7 \\times 10^{14}\\,L_{\\odot}$ for 1202$-$0725.  The inspection of the spectral energy distributions  in Fig.\\,~\\ref{sed_j1148} shows that in the UV/optical and  NIR/MIR (see below), where the emission is dominated by the  active nucleus, the SEDs of the high-$z$ sources match the template SEDs  from lower redshift AGN \\citep{rich06}  reasonably well. However, while the low-$z$ infrared SEDs show only one broad peak at short MIR wavelengths, we observe a second peak with considerable excess emission at FIR wavelengths in our targets.  Greybody fits to the data at $\\lambda > 160$\\,$\\mu$m yield a  temperature of $\\sim$\\,$50-60$\\,K for the FIR dust emission, which is likely powered by the vigorous star formation \\citep[see also][]{ber03,bee06,wan08}. These large contributions to the FIR emission,  which can presumably  be attributed to ongoing star formation suggesting rapid bulge build-up in the  host galaxies at high redshift, are missing in most lower redshift AGN. A  combination of the mean QSO SED and the single temperature greybody representing the contributions from star formation is able to match the observed photometry.  We note however that both objects presented here were selected as  SDP targets because of their known strong FIR/sub(mm) emission. In  principle, they could represent a small fraction of strongly star-forming objects,  while the majority of the high-$z$ QSOs may lack such a powerful FIR component.  An analogue for this situation may be found in the (mostly local) PG quasars \\citep{haas03}: when compared to the high-$z$ objects, many of the PG sources  show a similar mismatch in the FIR as seen for the SDSS QSO template, but a small  number of PG quasars reveal a comparable FIR excess as seen for the high-redshift QSOs.  Alternatively, stronger FIR emission from bulge build-up via star formation may be  more common at high $z$ than at low $z$, and we plan to explore these questions once data  for a greater number of high-redshift objects become available.  Both QSOs show more flux at NIR and MIR wavelengths than we would expect from the local templates. For 1202$-$0725 this may be understood to be due to the inclusion of the companion (which is at the same redshift as the quasar;  \\citealt{omo96a}, \\citealt{car02}) and which contributes roughly half of  the flux in many infrared and (sub)mm bands  \\citep[e.g.][]{omo96a,hin06,ion06}. For J1148+5251 the situation is somewhat different, because the QSO is clearly the dominating  source in flux at e.g. 8 and 24\\,$\\mu$m (observed, Fig.~\\ref{pic_j1148}). However, the exceptionally luminous  black hole in this QSO accretes close to its Eddington limit, which could possibly  lead to a larger fraction of dust being heated to high temperatures, resulting  in increased NIR emission.  For both objects a secondary component may contribute to the measured flux densities at FIR wavelengths, potentially leading to  an overestimate of the FIR luminosity of the QSO itself.  \\begin{table}[h] \\caption{Multi-wavelength data. All fluxes are given in mJy.} \\label{table1} \\begin{center} \\begin{tabular}{ccccc} \\hline\\hline $\\lambda_{\\rm obs}$ & SDSS J1148+5251        & refs & BR 1202$-$0725 & refs \\\\ in $\\mu$m          & $z=6.42$               &            & $z = 4.69$     & \\\\ \\hline 0.1450\\tablefootmark{a} & 0.0887          &  1  & 0.23$\\pm$0.02   & 11 \\\\ 1.2                     & 0.080$\\pm$0.004\\tablefootmark{b} &  1  & 0.32$\\pm$0.04\\tablefootmark{b}   & 12   \\\\ 1.6                     & 0.092$\\pm$0.004\\tablefootmark{b} &  2  & 0.50$\\pm$0.06\\tablefootmark{b}   & 12    \\\\ 2.2                     &  0.11$\\pm$0.005\\tablefootmark{b} &  2  & 0.41$\\pm$0.08\\tablefootmark{b}   & 12   \\\\ 3.6                     & 0.124$\\pm$0.002 &  3  & 0.461$\\pm$0.001 &  4    \\\\ 4.5                     & 0.140$\\pm$0.003 &  3  & 0.375$\\pm$0.001 &  4    \\\\ 5.8                     & 0.133$\\pm$0.010 &  3  & 0.410$\\pm$0.001 &  4    \\\\ 8.0                     & 0.241$\\pm$0.016 &  3  & 0.657$\\pm$0.002 &  4    \\\\ 16.0                     & 0.84           &  3  & --             &--      \\\\ 24.0                     & 1.520$\\pm$0.130 &  4  & 5.1$\\pm$0.6\\tablefootmark{c}     &  5   \\\\ 70.0                     & --             & -- & 15.0$\\pm$2.0\\tablefootmark{c}    &  5   \\\\ 100.0                     & 4.1$\\pm$1.0     &  5  & --             &--     \\\\ 160.0                     & 6.3$\\pm$2.0     &  5  & 39.8$\\pm$3.7\\tablefootmark{c} &  5   \\\\ 350.0                     &  21$\\pm$6       &  6  & 106$\\pm$7\\tablefootmark{c}       & 13     \\\\ 450.0                     & 24.7$\\pm$7.4    &  7  & 92$\\pm$38\\tablefootmark{c}       & 14     \\\\ 800.0                     & --             & -- & 50$\\pm$7\\tablefootmark{c}        & 14    \\\\ 850.0                     & 7.9$\\pm$0.7     &  7  & --            & --     \\\\ 900.0                     & --             & -- & 59$\\pm$11\\tablefootmark{c}       & 15     \\\\ 1100.0                    & --             & -- & 21$\\pm$5\\tablefootmark{c}        &14 \\\\ 1200.0                    & 5.0$\\pm$0.6     &  8  & --             &-- \\\\ 1330.0                    & 3.9$\\pm$0.8     &  9  & --             &-- \\\\ 1350.0                    & --             & -- & 16.0$\\pm$2.2\\tablefootmark{c}    & 16 \\\\ 2750.0                    & 0.28$\\pm$0.07   &  9  & --             &-- \\\\ 1.4\\,GHz                  & 0.055$\\pm$0.8   &  10 & 0.315$\\pm$0.080\\tablefootmark{c} & 17 \\\\ \\hline \\end{tabular} \\end{center} \\tablefoottext{a}{Refers to the 1450\\,\\AA~flux in the rest frame of the source.}\\\\ \\tablefoottext{b}{Calculated using the zero points presented in \\citet{coh03}.}\\\\ \\tablefoottext{c}{Total flux including the companion.} \\tablebib{(1)~\\citealt{fan03}; (2) \\citealt{iwa04}; (3) \\citealt{jia06}; (4) \\citealt{hin06}; (5) this work; (6) \\citealt{bee06}; (7) \\citealt{rob04}; (8) \\citealt{ber03}; (9) \\citealt{rie09}; (10) \\citealt{car04}; (11) \\citealt{sto96}; (12) \\citealt{skr06}; (13) \\citealt{ben99}; (14) \\citealt{isa94}; (15) \\citealt{ion06}; (16) \\citealt{omo96a}; (17) \\citealt{yun00}.} \\end{table}"
    },
    "1005/1005.4406_arXiv.txt": {
        "abstract": "Recently M\\'enard et al.\\ (hereafter MSFR) detected a subtle but systematic change in the mean color of quasars as a function of their projected separation from foreground galaxies, extending to comoving separations of $\\sim 10\\hMpc$, which they interpret as a signature of reddening by intergalactic dust. We present theoretical models of this remarkable observation, using smoothed particle hydrodynamic (SPH) cosmological simulations of a $(50\\hMpc)^3$ volume.  Our primary model uses a simulation with galactic winds and assumes that dust traces the intergalactic metals. The predicted galaxy-dust correlation function is similar in form to the galaxy-mass correlation function, and reproducing the MSFR data requires a dust-to-metal mass ratio of 0.24, about half the value in the Galactic ISM. Roughly half of the reddening arises in dust that is more than $100\\hKpc$ from the nearest massive galaxy. We also examine a simulation with no galactic winds, which predicts a much smaller fraction of intergalactic metals (3\\% vs.\\ 35\\%) and therefore requires an unphysical dust-to-metal ratio of 2.18 to reproduce the MSFR data.  In both models, the signal is dominated by sightlines with $E(g-i)=0.001-0.1$.  The no-wind simulation can be reconciled with the data if we also allow reddening to arise in galaxies up to several $\\times 10^{10} M_\\odot$.  The wind model predicts a mean visual extinction of $\\langle A_V \\rangle \\approx 0.0133$ mag out to $z=0.5$, with a sightline-to-sightline dispersion similar to the mean, which could be significant for future supernova cosmology studies.  Reproducing the MSFR results in these simulations requires that a large fraction of ISM dust survive its expulsion from galaxies and its residence in the intergalactic medium.  Future observational studies that provide higher precision and measure the dependence on galaxy type and environment will allow detailed tests for models of enriched galactic outflows and the survival of intergalactic dust. ",
        "introduction": "\\label{sec:intro} Heavy element enrichment in intergalactic gas has long been studied via absorption lines in quasar spectra and, in more restricted regimes, via the influence of metal species on the X-ray emission spectra of clusters and groups.  In the Galactic interstellar medium (ISM), roughly half of the mass of heavy elements is in solid phase --- interstellar dust --- which is thought to contain roughly 2/3 of interstellar carbon and the great majority of interstellar magnesium, silicon, and iron (\\citealt{Draine09} and references therein).  This solid-phase material is much more difficult to study in the intergalactic medium (IGM), in part because the expected level of dust extinction along any given line of sight is small, and in part because the continuum spectra of potential background sources --- quasars, galaxies, and supernovae --- show much greater variation than those of spectroscopically typed stars.  In a recent breakthrough study, M{\\'e}nard et al.\\ (2009, hereafter MSFR) used large samples of quasars and galaxies from the Sloan Digital Sky Survey (SDSS; \\citealt{York00}) to measure reddening by intergalactic dust, detecting a systematic change in quasar colors with increasing distance from foreground galaxies, measured to comoving projected separations of $\\sim 10\\hMpc$.  In this paper, we use hydrodynamic cosmological simulations to provide the first theoretical models of the MSFR results and to infer their implications for galactic outflows and the survival of dust in the IGM.  Ever since \\cite{Chandrasekhar52} modeled the opacity of interstellar matter using angular correlation statistics in the brightness fluctuations of the Milky Way, there have been efforts to constrain the presence of a diffuse dust component outside the Galaxy by measuring the extinction of background objects.  \\cite{Zwicky62} first claimed the existence of intracluster dust in Coma based on the extinction of light from background clusters, though his inferred level is certainly incompatible with more modern constraints. Quasars are powerful ``backlights'' for studying the transparency of the Universe across a wide range of cosmic time.  \\cite{OH84} modeled the effects of uniform and clumped dust distributions on the redshift evolution of quasar luminosity functions, but they concluded that observations did not clearly favor one specific dust distribution over another.  Taking advantage of the efficiency and accurate calibration afforded by CCD observations, \\cite{Zaritsky94} averaged hundreds of background galaxies to detect a mean color excess of $E(B-I)\\sim0.067$ behind two nearby galactic halos, arguing for the existence of dust out to at least $60\\hKpc$. \\cite{Menard08} also found evidence of dust in the halos of L$_*$ galaxies via studies of MgII absorbers. Confirming both findings, \\cite{Kaneda09} presented a direct detection of extended far-IR dust emission in the halo of NGC~253. The existence of dust in the intra{\\it cluster} medium has remained controversial, however, as reviewed by \\cite{Muller08}, who derived $E(B-R) = 0.005 \\pm 0.008$ from galaxies behind 458 RCS clusters.  Using a larger sample of $\\sim 10^4$ clusters from the SDSS, \\cite{Chelouche07} obtained significant detections of $E(g-i) \\approx\\,{\\rm few}\\times 10^{-3}$ out to several virial radii using the average colors of background quasars.  \\cite{Bovy08} inferred upper limits at roughly the same level for a sample of $z\\approx 0.05$ SDSS clusters using the spectra of early-type galaxies, though their results are not clearly inconsistent with Chelouche et al.'s.  MSFR instead examined the mean colors of 85,000 photometrically identified quasars \\citep{Richards04} as a function of angular separation from 20 million foreground galaxies in a 3800 square degree sky area from the SDSS.  They also used the quasar brightnesses to infer the galaxy-mass correlation function from weak lensing magnification.  They found that the galaxy-dust and galaxy-mass correlation functions track each other remarkably well, with a correlated dust surface density profile $\\Sigma_{\\rm dust} \\propto \\theta^{-0.8}$ detected out to angular separations that correspond to $\\sim 10\\hMpc$.  The color dependence of the reddening signal is consistent with ``standard'' interstellar dust, though the constraints are loose, implying $R_V \\equiv A_V/E(B-V) = 3.9 \\pm 2.6$.  The MSFR results provide direct evidence for extended intergalactic dust absorption, with an inferred total dust mass comparable to the amount of dust in galaxy disks.  Our primary model in this paper is based on a cosmological smoothed particle hydrodynamics (SPH) simulation that incorporates the momentum-driven galactic wind prescription of \\cite{OD06}, which successfully matches several observed aspects of gas-phase metal enrichment in galaxies and the IGM (see references in \\S\\ref{sec:sim}).  In momentum-driven wind models, radiation pressure on dust grains is the primary mechanism that sweeps gas out of galaxies \\citep{MQT05}.  We do not attempt to model the survival of dust in the IGM; rather, we treat the ratio of dust mass to heavy element mass as a free parameter to be inferred by matching the MSFR reddening measurements.  The metallicity of the outflowing gas is assumed to equal that of the star-forming gas, whose associated star formation effectively powers the outflow with radiation pressure.  The model here is almost the converse of the one treated by \\cite{Aguirre01}, who considered radiative expulsion of dust (without gas entrainment) and subsequent destruction as a source of IGM {\\it gas} enrichment.  As alternatives to the momentum-driven wind scenario, we also consider two models based on an SPH simulation without galactic winds.  The first assumes that the MSFR reddening signal arises in the simulation's enriched intergalactic gas, which comes mainly from ram pressure and tidal stripping in groups and clusters.  The second model allows additional absorption in low mass galaxies.  We hope to address a number of questions inspired by the MSFR observations. Can the simulation with galactic winds explain these observations?  What dust-to-metal mass ratio is required to do so, and what does this imply about the survival of dust in the IGM?  How is the dust distributed in the simulation?  Do the models without galactic winds provide a viable alternative explanation of the MSFR findings?  What further observations can test the model predictions and provide greater insight into the origin and evolution of intergalactic dust?  Section~\\ref{sec:sim} describes our simulations and presents visual maps of the predicted spatial distribution of intergalactic metals.  Section~\\ref{sec:res} presents our main results, including the predicted galaxy-dust correlations and projected reddening profiles, the contributions to the reddening signal from gas at different distances from high mass galaxies, the relative amounts of galactic and intergalactic dust, and the cumulative distribution functions of dust reddening. In section~\\ref{sec:dis}, we briefly discuss the issues of dust survival and extinction of high-redshift supernovae. Section~\\ref{sec:conc} summarizes our results with a look to potential future directions for observational studies. ",
        "conclusions": "\\label{sec:dis} For the Wind Model to succeed, we require that the dust-to-metal mass ratio in the IGM be comparable to that in the ISM, allowing only $\\sim 50\\%$ of the ISM dust to be destroyed during its expulsion from galaxies and subsequent residence in the IGM.  The validity of this assumption is by no means obvious, as the destruction timescales for $0.01\\,\\mu$m dust grains by thermal sputtering are $\\sim 10^{7.5}(n_H/10^{-3}\\cm^{-3})^{-1}$ years at $T=10^6\\K$ (\\citealt{DS79}, Figure 7), while wind particles in the simulation typically remain in the IGM for $\\sim 10^9$ years before reaccreting onto galaxies (\\citealt{Oppenheimer10}, Figure 2).  However, the sputtering rates decline rapidly towards lower temperatures (e.g., a factor of 300 lower at $T=10^5\\K$), and with the wind implementation used in this simulation most ejected gas never rises above a few $\\times 10^4\\K$.  UV or X-ray background photons are another possible destruction mechanism for IGM dust, but the intergalactic radiation field is much lower intensity than the radiation field dust grains already encounter in galactic star-forming regions.  A detailed consideration of dust survival in the IGM is beyond the scope of this initial investigation, but the MSFR results clearly raise it as an important subject for further study.  The combination of their measurements with our models gives a fairly clear idea of what is required: survival of a substantial fraction of ejected dust, and an extinction curve that has roughly the color dependence of ISM dust.  The temperature sensitivity of thermal sputtering could lead to preferential destruction of ejected dust in the higher mass halos that host a shock heated gas halo (see \\citealt{Birnboim03,Keres05,Dekel06,Keres09a}).  In the Wind Model, most wind particles in halos with $M < 10^{13} M_\\odot$ have $T < 10^5\\K$, but about 2/3 of the wind particles in halos with $M>10^{13} M_\\odot$ have $T > 3\\times 10^6\\K$.  If sputtering does destroy intergalactic dust at these temperatures, it could produce distinctive drops in the galaxy-reddening correlation when it is evaluated for massive galaxies or for galaxies in dense environments.  The recent study of \\cite{McGee10}, which examines the correlation of background quasar colors with projected separation from galaxy groups of varying richness, provides some hint of such an effect, but their innermost point is at $r = 1\\hMpc$, close to the virial radius of typical group mass halos. Moreover, the \\cite{Chelouche07} measurements provide direct evidence for dust survival in the cluster IGM.  \\cite{Draine09} argues that ISM dust must form largely {\\it in situ}, from the depletion of gas phase metals onto seed grains from supernovae and ejected stellar envelopes, and that the dust abundance is determined by an equilibrium between growth and destruction mechanisms. Growth rates would be much slower in the low density IGM, but 2-body destruction processes would have the same density scaling, so a similar equilibrium abundance could arise. \\begin{figure} \\centering \\resizebox{0.85\\columnwidth}{!}{\\includegraphics{f7.eps}} \\caption{Comparison of reddening maps in the three dust models, each normalized to the MSFR result at 1$\\hMpc$.  The regions shown are a quadrant of the simulation cube, $25\\hMpc\\times25\\hMpc\\times50\\hMpc$.  Red/blue regions indicate heavily/lightly reddened fields through the simulation, with the color scale running from $E(g-i)=10^{-5}$ up to $0.1$, logarithmically.} \\label{fig:extmap} \\end{figure}  Intergalactic dust could have an important impact on supernova cosmology studies, dimming and reddening supernovae increasingly with redshift.  Given perfect knowledge of supernova intrinsic colors and the shape of the extinction curve, intergalactic extinction would be corrected automatically along with extinction in the supernova host galaxy.  However, the intergalactic and host galaxy dust components could have different extinction curves, complicating the analyses.  Furthermore, the presence of intergalactic dust means that no high-redshift supernovae have zero extinction, contrary to the assumption usually made in global models of the supernova population.  MSFR give a rough estimate of the average extinction implied by their results, $\\langle A_V \\rangle = 0.03$ mag at $z=0.5$, but the inference of a mean extinction from the {\\it correlated} reddening signal is (as MSFR emphasize) sensitive to uncertain choices of radial profile cutoffs and luminosity extrapolations. (See \\citealt{Menard09b} for further discussion.)  Our simulation provides an explicit physical model of the IGM metal distribution, and once the dust-to-metal mass ratio is set by matching the MSFR data, it is straightforward to compute the mean extinction by intergalactic dust.  In the Wind Model, the comoving dust density $\\rho_{dust}=6.424\\times10^{-35}$g\\,cm$^{-3}$. Assuming constant comoving dust density and extinction law $A_\\lambda\\propto\\lambda^{-1.2}$, we can compute mean extinction to any given redshift $z$ as \\begin{equation} \\langle A_V \\rangle = \\frac{2.5\\,K_{\\rm ext}\\,c}{{\\rm ln}(10)} \\int_0^z{(1+z)^{3.2}H(z)^{-1}\\rho_{\\rm dust}{\\rm d}z}, \\end{equation} where $H(z)$ is the Hubble constant at redshift $z$, respectively. The $(1+z)^{3.2}$ term arises from the combination of cosmic areal expansion and rest-frame to observed V-band extinction conversion. This yields a predicted mean extinction $\\langle A_V \\rangle = 0.0133$ mag and a sightline-to-sightline variance of $\\Delta^2 = (0.0272\\,{\\rm mag})^2$ for supernovae at $z=0.5$. More accurate estimates and an extrapolation to higher redshifts will require the combined analysis of many simulation outputs to track evolution.  Our estimated extinction is small compared to the statistical and systematic errors of current surveys, but it is comparable to the ambitious goals set for next-generation surveys of several thousand high-redshift supernovae.  At higher redshifts, furthermore, the (stronger) near-UV dust extinction gets redshifted into the observed-frame optical bands.  We have assumed SMC dust for this calculation; since we normalize to MSFR's $E(g-i)$ measurements, the important assumption is the value of $A_V/E(g-i) = 4.7$."
    },
    "1005/1005.3821_arXiv.txt": {
        "abstract": "{The {\\it Herschel} Lensing Survey  (HLS) takes advantage of gravitational lensing by massive galaxy clusters to sample a population of high-redshift galaxies which are too faint to be detected above the confusion limit of current far-infrared/submillimeter telescopes. Measurements from 100--500\\,$\\mu$m bracket the peaks of the far-infrared spectral energy distributions of these galaxies, characterizing their infrared luminosities and star formation rates.  We introduce initial results from our science demonstration phase observations, directed toward the Bullet cluster (1E0657-56). By combining our observations with LABOCA 870\\,$\\mu$m and AzTEC 1.1\\,mm data we fully constrain the spectral energy distributions of 19 MIPS 24\\,$\\mu$m-selected galaxies which are located behind the cluster.  We find that their colors are best fit using templates based on local galaxies with systematically lower infrared luminosities.This suggests that our sources are not like local ultra-luminous infrared galaxies in which vigorous star formation is contained in a compact highly dust-obscured region. Instead, they appear to be scaled up versions of lower luminosity local galaxies with star formation occurring on larger physical scales.} ",
        "introduction": "\\begin{figure*}[ht] \\centering \\includegraphics[width=5.0in]{f1_14693.eps} \\caption{The most highly magnified galaxy in our sample, HLS18, is the same LIRG first detected by \\citet{Wilson2008} with AzTEC at 1.1\\,mm. {\\it Left panel}: 1.5$\\arcmin \\times $2$\\arcmin$ thumbnails of the source at a variety of wavelengths highlight the importance of multi-wavelength data to guide the interpretation of the lower resolution maps.  IRAC data reveal three lensed images of the galaxy indicated with white circles, although only emission from the two brightest images is detected in the {\\it Herschel} maps. {\\it Right panel}: \\citet{Rieke2009} SED template fit to the source.  The fluxes include the contribution from both of the brightest images of the galaxy, indicated by the two left-most circles in the thumbnails.  The data points have not been de-magnified. The 24\\,$\\mu$m flux is plotted for reference, although it was not used in the fit.  The best-fit modified blackbody is shown with a dotted line and is consistent with previous estimates \\citep{Wilson2008,Gonzalez2009,Rex2009}. } \\label{lssrc} \\end{figure*} Luminous (and ultra-luminous) infrared galaxies ((U)LIRGS) are a population of massive star-forming galaxies which contain significant amounts of dust, absorbing the ultraviolet emission from newly formed stars within them and re-radiating this energy at far-infrared (FIR) wavelengths. The FIR emission of these sources is therefore strongly correlated with their level of star formation activity \\citep[e.g.][]{Kennicutt1998}.  {\\it {Spitzer}} MIPS measurements from 24--160\\,$\\mu$m bracket the peak of the dust emission from nearby (U)LIRGS.  Therefore, MIPS observations have been very successful in characterizing the spectral energy distributions (SEDs) and star formation rates (SFRs) of the local population \\citep[e.g.][]{Rieke2009}.  Similar measurements at FIR/submillimeter (submm) wavelengths are required to constrain the SEDs of their higher redshift counterparts.  Together the PACS \\citep{Poglitsch2010} and SPIRE \\citep{Griffin2010} instruments on-board {\\it Herschel} \\citep{Pilbratt2010} provide sensitive measurements of this population at five wavelengths from 100--500\\,$\\mu$m, sampling the rest-frame peak of the FIR emission out to a redshift of $z = 4$.  The depth of submm maps is ultimately limited by confusion noise resulting from a high density of sources relative to the angular resolution of the telescope.  Gravitational lensing by massive galaxy clusters provides the only means to sample an abundant population of intrinsically faint or high-redshift infrared (IR) galaxies which lie below this limit. The {\\it Herschel} Lensing Survey (HLS) \\citep{Egami2010} will target $\\sim$ 40 galaxy clusters with PACS and SPIRE to compile the first significant submm catalog of these galaxies. In this letter we present initial results from our science demonstration phase observations: 8 $\\arcmin \\times$ 8$\\arcmin$ PACS maps and 17 $\\arcmin \\times$ 17$\\arcmin$ SPIRE maps centered on the Bullet cluster (1E0657-56).  Details of these observations are presented in \\citet{Egami2010}.  This is a unique target because of its strong lensing potential and the richness of existing ancillary data.  We present the FIR properties of 19 MIPS 24\\,$\\mu$m-selected galaxies which are located behind the Bullet cluster.  We compare the shapes of their SEDs with those of local (U)LIRGS, and compare their measured IR luminosities with predictions which extrapolate this quantity from the observed MIPS 24\\,$\\mu$m flux. Subsequent uses of MIPS in the text refer to the MIPS 24\\,$\\mu$m band. ",
        "conclusions": ""
    },
    "1005/1005.1223_arXiv.txt": {
        "abstract": "In the {\\it IRAS\\/} [12]--[25],  [25]--[60] colour--colour diagram, RV Tauri stars are found to populate cooler temperature regions $(T<600\\,\\rmn{K})$, distinctly different from those occupied by  the oxygen and carbon Miras. The {\\it IRAS\\/} fluxes are  consistent with the dust density in the envelope varying as the inverse square of the radial distance, implying that the grain formation processes in these objects are most probably continuous and not sporadic. It is found that the spectroscopic subgroups A and B are well separated in the far-infrared two-colour diagram, with group B objects having systematically cooler dust envelopes. We interpret this as being due to a difference in the nature of grains, including the chemical composition, in the two cases. ",
        "introduction": "It has been well established that RV Tauri variables  possess infrared emission far in excess of their expected  blackbody continuum, arising from their extended cool dust envelopes \\citep{b7,b5,b6}. Recently, \\citet{b15} and \\citet{b9} have given detailed descriptions of the near-infrared properties of RV Tauri stars. In this paper we present an analysis of the {\\it IRAS\\/} data of RV Tauri stars with the help of the far-infrared two-colour diagram and a grid computed using a simple model of the dust envelope. Such two-colour plots have already been employed extensively by several investigators to study the circumstellar envelopes around oxygen-rich and carbon-rich objects which are in the late stages of stellar evolution \\citep{b10,b25,b23,b24}.  Table 1 summarizes the basic data on the 17 objects  detected at \\hbox{60\\,$\\umu$m}. Apart from the {\\it IRAS\\/} identification and the flux densities at 12-, 25-, 60-  and 100-$\\umu$m wavebands, it gives the spectroscopic groups of \\citet{b20}, the light-curve classes of \\citet{b13} and the periods of light variation. The list, which contains  about 20 per cent of all the known RV Tauri stars, is  essentially the same as that given by \\citet{b12}. The spectroscopic subgroups are from either \\citet{b20} or \\citet{b15}. \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Data on the RV Tauri stars detected by {\\it IRAS}.} \\begin{tabular}{@{}llrrrrlrlr@{}} \\hline Name     &            & \\multicolumn{4}{c}{Flux density (Jy)% \\footnote{Observed by {\\em IRAS}.}}\\\\ Variable & {\\it IRAS} & 12$\\,\\umu$m & 25$\\,\\umu$m & 60$\\,\\umu$m & 100$\\,\\umu$m & Sp. & Period & Light- & $T_0\\,(\\rmn{K})$ \\\\ &  &  &  &  &  & group & (d) & curve \\\\ &  &  &  &  &  &       &     & type  \\\\ \\hline TW Cam & 04166$+$5719 & 8.27 & 5.62 & 1.82 & $<$1.73 & A & 85.6 & a & 555 \\\\ RV Tau & 04440$+$2605 & 22.53 & 18.08 & 6.40 & 2.52 & A & 78.9 & b & 460 \\\\ DY Ori & 06034$+$1354 & 12.44 & 14.93 & 4.12 & $<$11.22 & B & 60.3 &  & 295 \\\\ CT Ori & 06072$+$0953 & 6.16 & 5.57 & 1.22 & $<$1.54 & B & 135.6 &  & 330 \\\\ SU Gem & 06108$+$2734 & 7.90 & 5.69 & 2.16 & $<$11.66 & A & 50.1 & b & 575 \\\\ UY CMa & 06160$-$1701 & 3.51 & 2.48 & 0.57 & $<$1.00 & B & 113.9 & a & 420 \\\\ U Mon  & 07284$-$0940 & 124.30 & 88.43 & 26.28 & 9.24 & A & 92.3 & b & 480 \\\\ AR Pup & 08011$-$3627 & 131.33 & 94.32 & 25.81 & 11.65 & B & 75.0 & b & 450 \\\\ IW Car & 09256$-$6324 & 101/06 & 96.24 & 34.19 & 13.07 & B & 67.5 & b & 395 \\\\ GK Car & 11118$-$5726 & 2.87 & 2.48 & 0.78 & $<$12.13 & B & 55.6 &  & 405 \\\\ RU Cen & 12067$-$4508 & 5.36 & 11.02 & 5.57 & 2.01 & B & 64.7 &  & 255 \\\\ SX Cen & 12185$-$4856 & 5.95 & 3.62 & 1.09 & $<$1.50 & B & 32.9 & b & 590 \\\\ AI Sco & 17530$-$3348 & 17.68 & 11.46 & 2.88 & $<$45.62 & A & 71.0 & b & 480 \\\\ AC Her & 18281$+$2149 & 41.47 & 65.33 & 21.12 & 7.79 & B & 75.5 & a & 260 \\\\ R Sct  & 18448$-$0545 & 20.88 & 9.30 & 8.10 & $<$138.78 & A & 140.2 & a \\\\ R Sge  & 20117$+$1634 & 10.63 & 7.57 & 2.10 & $<$1.66 & A & 70.6 & b & 455 \\\\ V Vul  & 20343$+$2625 & 12.39 & 5.72 & 1.29 & $<$6.96 & A & 75.7 & a & 690\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}  \\section[]{Description of the Envelope\\\\* Model}  If we assume that the dust grains in the envelope are  predominantly of the same kind and are in thermal  equilibrium, the luminosity at frequency $\\nu$ in the infrared is given by \\begin{equation} L(\\nu)=\\mskip-12mu\\int\\limits_{\\rmn{envelope}}\\mskip-12mu \\rho(r)Q_{\\rmn{abs}}(\\nu)B[\\nu,T_{\\rmn{g}}(r)]\\exp [-\\tau(\\nu,r)]\\> \\rmn{d}V, \\end{equation} where $Q_{\\rmn{abs}}(\\nu)$ is the absorption efficiency at frequency $\\nu$, $\\rho(r)$            is the dust grain density, $T_{\\rmn{g}}(\\nu)$    is the grain temperature, $B[\\nu,T_{\\rmn{g}}(r)]$  is the Planck function, and $\\tau(\\nu,r)$        is the optical depth at distance {\\it r\\/} from the centre of the star.  The temperature $T_{\\rmn{g}}(r)$ is determined by the condition of energy balance: amount of energy radiated = amount of energy absorbed. The amount of energy absorbed at any point is proportional to the total available energy at that point, which consists of: \\begin{enumerate} \\item the attenuated and diluted stellar radiation; \\item scattered radiation, and \\item reradiation from other grains. \\end{enumerate}  Detailed solutions of radiative transfer in circumstellar  dust shells by \\citet{b21,b22} indicate that the effect of heating by other grains becomes significant only at large optical depths at the absorbing frequencies $[\\tau(\\rmn{UV})\\gg 10]$, and at optical depths $\\tau(\\rmn{UV})<1$ the grains have approximately the same temperature that they would have if they were seeing the starlight unattenuated and no other radiation.  The Planck mean optical depths of circumstellar envelopes  around several RV Tauri stars, derived from the ratios of the luminosities of the dust shell (at infrared wavelengths) and the star, range from 0.07 to 0.63 \\citep{b9}. There is much uncertainty in the nature of the optical properties of dust grains in the envelope. The carbon-rich RV Tauri stars are also reported to show the 10-$\\umu$m silicate emission feature typical of oxygen-rich objects \\citep{b6,b19}. The pure terrestrial silicates or lunar silicates are found to be completely unsuitable to account for the infrared emission from circumstellar dust shells around M-type stars \\citep{b21}. We assume that the absorption efficiency $Q_{\\rmn{abs}} (\\nu)$ in the infrared varies as $\\nu^{\\gamma}$. ${\\gamma}=1$ appears to provide a reasonable fit in a variety of sources \\citep*{b11,b12}. Under these circumstances the condition of energy balance implies that the dust temperature $T_{\\rmn{g}}$ will vary as $r^{\\beta}$.  In view of the low value of the observed Planck mean  optical depth for the stellar radiation and the nature of the assumed frequency dependence of the absorption efficiency, the extinction of the infrared radiation by the dust envelope can be neglected. If we consider the envelope to be spherically symmetric, equation (1) reduces to \\begin{equation} L(\\nu)=\\!\\!\\int_{r_{1}}^{r_{2}}\\!\\!4\\upi r^2\\rho(r)\\> Q_{\\rmn{abs}}(\\nu)B[\\nu,T_{\\rmn{g}}(r)]\\> {\\rmn{d}}r, \\end{equation} where $r_1$ and $r_2$ are the inner and outer radii of the shell. For a dusty density distribution $\\rho(r)\\propto r^{\\alpha}$ and $r_2\\gg r_1$, equation (2) reduces to \\begin{equation} L(\\nu)\\propto \\nu^{2+\\gamma-Q}\\int_{X_0}^{\\infty}{{x^Q}\\over {\\rmn{e}^x-1}}\\rmn{d}x , \\end{equation} where $Q=-(\\alpha+\\beta+3)/\\beta$ and $X_0=(h\\nu /kT_0)$. $T_0$ represents the temperature at the inner boundary of the dust shell where grains start condensing. In a steady radiation pressure driven mass outflow in the optically thin case, values of $\\alpha$ lie near $-2$ \\citep{b8}. $\\gamma$ and $\\beta$ are related by $\\beta=-2/(\\gamma+4)$.  In the {\\it IRAS\\/} Point Source Catalog \\citep[PSC,][]{b2}, the flux densities have been quoted at the effective wavelengths 12, 25, 60 and \\hbox{100\\,$\\umu$m}, assuming a flat energy spectrum $[\\nu F(\\nu)=1]$ for the observed sources. For each model given by equation (3), using the relative system response, the colour-correction factors \\citep{b3} in each of the {\\it IRAS\\/} passbands were calculated and the fluxes were converted into flux densities expected for a flat energy distribution, as assumed in the {\\it IRAS\\/} PSC, so that the computed colours can be directly compared with the colours determined from the catalogue quantities. Such a procedure is more appropriate than correcting the {\\it IRAS\\/} colours for the energy distribution given by a particular model and then comparing them with those computed by the model.\\footnote{An example of a footnote.}  \\subsection{Colour--colour diagram}  The IR colour is defined as \\[ [\\nu_1]-[\\nu_2]=-2.5\\log [f(\\nu_1)/f(\\nu_2)], \\] where $\\nu_1$ and $\\nu_2$ are any two wavebands and $f(\\nu_1)$ and $f(\\nu_2)$ are the corresponding flux  densities assuming a flat energy spectrum for the source. In Fig.~1, we have plotted the [25]--[60] colours  of RV Tauri stars against their corresponding [12]--[25]  colours derived from the {\\it IRAS\\/} data. Filled circles  represent stars of group A and open circles stars of group B. The two sets of near-parallel lines represent the loci of constant inner shell temperature $T_0$ and the quantity $Q$ defined above. The models correspond to the case of absorption efficiency $Q_{\\rmn{abs}}(\\nu)$ varying as $\\nu$ (with $\\gamma=1$ and hence $\\beta=-0.4$). We have omitted R Sct in Fig.~1 because it shows a large deviation from the average relation shown by all the other objects. R Sct has a comparatively large excess at 60$\\,\\umu$m, but the extent of a possible contamination by the infrared cirrus \\citep{b16} is unknown. \\citet{b9} found no evidence of the presence of a dust envelope at near-IR wavelengths and the spectrum was consistent with a stellar continuum. This explains why R Sct lies well below the mean relation shown by stars of groups A and C between the [3.6]--[11.3] colour excess and the photometrically determined (Fe/H) \\citep{b4}. R Sct has the longest period of 140$\\,$d among the RV Tauri stars detected at far-infrared wavelengths and does not have the 10-$\\umu$m emission feature seen in other objects \\citep{b5,b19}. R Sct is probably the most irregular RV Tauri star known \\citep{b17}.  \\begin{figure} \\vspace{302pt} \\caption{Plot of [25]--[60] colours of RV  Tauri stars against their [12]--[25] colours after  normalizing as indicated in \\citet{b3}. Some of the objects are identified by their variable-star names. Typical error bars are shown in the bottom right-hand corner. The lines represent the loci for constant inner shell temperature and the quantity $Q$. Note the separation of group A and B stars at $T_0 \\sim$ 460$\\,$\\,K. Positions occupied by a sample of carbon and oxygen Miras are also shown. The $Q=1.0$ line differs from the blackbody line by a maximum of $\\sim 0.05$.} \\end{figure} The inner shell temperatures $(T_0)$ derived for the various objects are also given in Table~1 and we find the majority of them to have temperatures in the narrow range 400--600$\\,$K. If the dependences of $Q_{\\rmn{abs}}(\\nu)$ on $\\nu$ and $\\rho(r)$ on $r$ are similar in all the objects considered, then in the colour--colour diagram they all should lie along a line corresponding to different values of $T_0$ and in Fig.~1 we find that this is essentially the  case. In view of the quoted uncertainties in the flux measurements, we cannot attach much significance to the scatter in Fig.~1.  At \\hbox{100\\,$\\umu$m} the infrared sky is characterized by emission, called infrared cirrus, from interstellar dust on all spatial scales \\citep{b16}, thereby impairing the measurements at far-infrared wavelengths. In Fig.~2, we have plotted the [60]--[100] colours of the six RV Tauri stars detected at \\hbox{100\\,$\\umu$m} against their [25]--[60] colours, along with the grid showing the regions of different values for inner shell temperature $T_0$ and the quantity $Q$, as in Fig.~1. The results indicated by Fig.~2 are consistent with those derived from Fig.~1. AR Pup shows a large excess at \\hbox{100\\,$\\umu$m} but, in view of the large values for the cirrus flags given in the catalogue, the intrinsic flux at \\hbox{100\\,$\\umu$m} is uncertain.  \\subsection{Radial distribution of dust}  \\begin{figure*} \\vspace*{174pt} \\caption{Plot of the [60]--[100] colours of RV Tauri stars against their [25]--[60] colours after normalizing as indicated in \\citet{b3}. The solid lines represent the loci for constant inner shell temperature and the quantity $Q$. The dashed line shows the locus for a blackbody distribution.} \\end{figure*}  From Fig.~1, it is evident that all RV Tauri stars lie between the lines corresponding to $Q=1.5$ and 0.5. With \\[ \\alpha=-(1+Q)\\beta-3, \\] these values suggest limits of $r^{-2.0}$ and $r^{-2.4}$ for the dust density variation, indicating a near-constant mass-loss rate. \\citet{b12} has suggested that the density in the circumstellar envelope around RV Tauri stars varies as $r^{-1}$, implying a mass-loss rate that was greater in the past than it is currently. By fitting a power law to the observed fluxes, such that $f(\\nu)$ varies as $\\nu^q$, values of $q$ determined by him for the various objects given in Table~1 lie in the range 0.6--1.2, with a mean $\\skew5\\bar q=0.98$. The assumption of a power law corresponds to the case of $X_0=0$ in equation (3) and hence we get \\[ q=2+\\gamma -Q. \\] Since we assume that $Q_{\\rmn{abs}}(\\nu)$ varies as $\\nu$, the resulting value for $Q$=2.0. None of the objects is found to lie in the corresponding region in the colour--colour diagram. Even this extreme value for $Q$ implies a density which varies as $r^{-1.8}$.  \\citet{b9} have reported that the simultaneous optical and near-IR data of AC Her can be fitted by a combination of two blackbodies at 5680 and 1800\\,K, representing, respectively, the stellar and dust shell temperatures, and suggested that in RV Tauri stars the grain formation is a sporadic phenomenon and not a continuous process. Apparently, they have been influenced by the remark by \\citet{b7} that their data in the 3.5--11$\\,\\umu$m region of AC Her indicated a dust temperature of $\\sim$300\\,K. We find that the {\\it K--L\\/} colours given by \\citet{b5}, \\citet{b15} and \\citet{b9} are all consistent with each other. Surely, hot dust ($\\sim 1800\\,$K), if present at the time of observations by \\citet{b9}, would have affected the {\\it K--L\\/} colour significantly. AC Her, like other members of its class, is found to execute elongated loops in the ({\\it U--B\\/}), ({\\it B--V\\/}) plane \\citep{b20}, indicating that significant departure of the stellar continuum from the blackbody is to be expected. Further, their data show only a marginal excess at the near-IR wavelengths. We feel that the case for the existence of hot dust around AC Her and hence for the sporadic grain formation around RV Tauri stars is not strong. In Fig.~3 we find that AC Her and RU Cen lie very close to R Sct which, according to \\citet{b9}, shows no evidence for the presence of a hot dust envelope.  \\subsubsection{Comparison with oxygen and carbon Miras}  In Fig.~1 we have also shown the positions of a sample of oxygen-rich and carbon-rich Miras. At the low temperatures characteristic of the Miras, a part of the emission at 12$\\,\\umu$m comes from the photosphere. For a blackbody at 2000$\\,$K, the ratio of fluxes at wavelengths of 12 and 2$\\,\\umu$m $(f_{12}/f_{2})\\sim 0.18$. The Miras shown in Fig.~1 have $(f_{12}/f_{2})$ ratios larger than twice the above value. It is clear that the three groups of objects populate three different regions of the diagram. \\citet{b10} have already noticed that there are distinct differences between the {\\it IRAS\\/} colours of oxygen-rich and carbon-rich objects. On the basis of an analysis, using a bigger sample of bright giant stars in the {\\it IRAS\\/} catalogue, this has been interpreted by \\citet{b25} as being due to a systematic difference in the dust grain emissivity index. U Mon shows the 10-$\\umu$m silicate emission convincingly and, in most of the other objects for which low-resolution spectra in the near-infrared have been reported \\citep{b5,b19}, the 10-$\\umu$m emission may be partly attributed to silicates. Hence it is reasonable to expect that, in the envelopes around at least some of the RV Tauri stars, the dust grains are predominantly of silicates, as in the case of oxygen Miras \\citep{b21}. The fact that none of the RV Tauri stars is found in the region of the two-colour diagram occupied by the oxygen Miras indicates that the emissivity indices of the silicate grains in the two cases are different. Because of the higher temperatures and luminosities, the environment of grain formation will be different in RV Tauri stars.  \\subsubsection{Correlation with subgroups}  \\citet{b20} have identified three spectroscopic subgroups, which are designated as groups A, B and C. Objects of group A are metal-rich; group C are metal-poor; group B objects are also metal-poor, but  show carbon enhancements \\citep{b20,b14,b4,b1}. It is interesting to see that Table~1 contains no group C objects and that in Fig.~1 there is a clear separation of the two spectroscopic subgroups A and B, with the demarcation  occurring at an inner shell temperature of about 450$\\,$K, group B stars having lower temperatures than group A. SX Cen is the only exception. \\citet{b14} has reported that metal lines are stronger in SX Cen than in other group B objects. It may be worth noting that SX Cen has the shortest period among the 100 or so objects with the RV Tauri classification. RU Cen has the coolest inner shell temperature, as already suggested by the near-infrared spectrum \\citep{b6}. \\begin{figure} \\vspace*{174pt} \\caption{Plot of ({\\it K--L\\/}) colours of RV Tauri stars detected by {\\it IRAS\\/} against their corresponding ({\\it J--K\\/}) colours. The position of AR Pup is indicated. The three objects lying close to the blackbody line are AC Her, RU Cen and R Sct.} \\end{figure}  Group B objects follow a different mean relationship from those of group A, having systematically larger 11-$\\umu$m excess for a given excess at 3$\\,\\umu$m \\citep{b15}. For a general sample of RV Tauri stars, the distinction between the oxygen-rich and carbon-rich objects is not that apparent in the {\\it JHKL\\/} bands. In Fig.~3 we have plotted the near-IR magnitudes of the objects given in Table~1 (except V Vul which has no available measurements) in the {\\it J--K, K--L\\/} plane. The colours,  taken from \\citet{b15} and \\citet{b9}, are averaged if more than one observation exists, because the internal agreements are found to be often of the order of observational uncertainties, in accordance with the earlier finding by \\citet{b5} that variability has relatively little effect on colours. Barring RU Cen and AC Her, it is evident that stars belonging to group B show systematically larger excesses at {\\it L\\/} band for a given excess at {\\it K}. The low excesses at near-IR wavelengths for AC Her and RU Cen are consistent with the very low dust temperatures indicated by the far-infrared colours. \\begin{figure*} \\vbox to 220mm{\\vfil Landscape figure to go here. This figure was not part of the original paper and is inserted here for illustrative purposes.\\\\ See the author guide for details (section 2.2 of \\verb|mn2eguide.tex|) on how to handle landscape figures or tables. \\caption{} \\vfil} \\label{landfig} \\end{figure*}  It is already well established that from {\\it UBV\\/} photometry one can distinguish between groups A and B,  members of group A being significantly redder than those of group B \\citep{b20}. Similarly, \\citet{b4} has found that the two spectroscopic groups are well separated in the DDO colour--colour diagrams when mean colours are used for the individual objects.  The clear separation of the spectroscopic subgroups A and  B in the IR two-colour diagram suggests that the natures of dust grains in the envelopes in the two cases are not  identical. This is to be expected because of the differences in the physical properties of the stars themselves. The average colours of group B stars are bluer than group A, but the envelope dust temperatures of B are cooler than those of A. The near-IR spectra of AC Her and RU Cen are extremely similar \\citep{b6}. The striking similarities in the optical spectra of AC Her and RU Cen have been pointed out by Bidelman \\citep{b18}. We feel that the physical properties, including the chemical composition, of the grains  formed in the circumstellar envelope strongly depend on those of the embedded star. This, probably, explains the diversity of the energy distributions of RV Tauri stars in the near-infrared found by \\citet{b6}. On the basis of the observed differences in chemical abundances and space distribution of RV Tauri stars, \\citet{b15} has already pointed out that there is no direct evolutionary connection between group A and group B objects, thus ruling out the possibility that group B objects are the evolutionary successors of group A, in which grain formation has stopped and the cooler temperatures for the former are caused by an envelope expansion.  \\citet{b13} have subdivided RV Tauri stars into two classes, RVa and RVb, on the basis of their light curves; the former shows a constant mean  brightness, whereas the latter shows a cyclically varying  mean brightness. Extensive observations in the near-infrared show that, on average, RVb stars are redder than RVa stars, and \\citet{b15} has suggested that in RVb stars dust shells are denser in the inner regions and hence radiate strongly in the 1--3$\\,\\umu$m region. Fig.~3 confirms this; RVb objects show systematically larger ({\\it J--K\\/}) and ({\\it K--L\\/}) colours than RVa objects. Apparently, there is no distinction between objects of the two light-curve types at far-infrared wavelengths (Fig.~1). ",
        "conclusions": "In the [12]--[25], [25]--[60] colour diagram, RV Tauri stars populate cooler temperature regions $(T<600 \\,\\rmn{K})$, distinctly different from those occupied by the oxygen and carbon Miras. Using a simple model in which \\begin{enumerate} \\item the envelope is spherically symmetric, \\item the IR-emitting grains are predominantly of the same kind, and \\item in the infrared the absorption efficiency $Q_{\\rmn{abs}} (\\nu)\\propto\\nu$, \\end{enumerate} we find that the {\\it IRAS\\/} fluxes are consistent with the density in the envelope $\\rho(r)\\propto r^{-2}$, where {\\it r\\/} is the radial distance. Such a dependence for the dust density implies that the mass-loss rates in RV Tauri stars have not reduced considerably during the recent past, contrary to the suggestion by \\citet{b12}. In the two-colour diagram, the blackbody line and the line corresponding to $\\rho(r)\\propto r^{-2.2}$ nearly overlap and the present data are insufficient to resolve between the two cases. The latter case is more physically reasonable, however.  The spectroscopic subgroups A and B are well separated in  the {\\it IRAS\\/} two-colour diagram, with group B objects  having systematically cooler dust envelopes. If we consider only the objects detected by {\\it IRAS}, we find that stars belonging to group B show systematically larger excess at {\\it L\\/} band for a given excess at {\\it K}. Apparently, there is no correlation between the light-curve types (RVa and RVb) and the far-infrared behaviour of these objects. It is fairly certain that the physical properties, including the chemical composition, of the embedded stars are directly reflected by those of the dust grains. Most probably, the grain formation process in RV Tauri stars is continuous and not sporadic as suggested by \\citet{b9}."
    },
    "1005/1005.4589_arXiv.txt": {
        "abstract": "Using a bistatic radar echo sounding (RES) system developed for calibration of the RICE particle astrophysics experiment at the South Pole, we have studied radio frequency (RF) reflections off the bedrock. The total propagation time of $\\sim$ns-duration, vertically (${\\hat z}$) broadcast radio signals, as a function of polarization orientation in the horizontal plane, provides a direct probe of the geometry-dependence of the ice permittivity to a depth of 2.8 km. We observe clear birefringent asymmetries along ${\\hat z}$ in the lowest half of the ice sheet, at a fractional level $\\sim$0.3\\%. % This result is in contrast to expectations based on measurements at Dome Fuji, for which birefringence was observed in the upper 1.5 km of the ice sheet. This effect, combined with the increased radio frequency attenuation expected near the bedrock, renders the lower half thickness of South Polar ice less favorable than the upper half of the ice sheet in terms of its ultra-high energy neutrino detection potential. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.3360_arXiv.txt": {
        "abstract": "{ We report the results of multifrequency-VLBI observations of GHz-Peaked-Spectrum (GPS) radio sources. The VLBI structure and component spectra of some GPS sources are presented. Our VLBI results show that about 80\\% of the GPS galaxies exhibit a compact double or CSO-like structure, while the GPS quasars tend to show a core-jet. The component spectra of the GPS galaxies are often steep/convex, and the core has a flat spectrum but it is usually hidden or weak. In addition, we studied the variability of GPS sources by comparing new flux density measures, acquired with the Urumqi 25m telescope at 4.85 GHz, with previous 87GB data. The results show that 44\\% of the GPS quasars varied higher than 10\\% in passed 20 years, while the fraction is only 12\\% for the GPS galaxies meaning that the GPS quasars are much more variable than GPS galaxies. In total, 25\\% of GPS sources show $>10\\%$ variability at 4.85 GHz in our sample.} ",
        "introduction": "GHz-Peaked-Spectrum (GPS) radio sources are powerful ($P_{\\rm1.4~GHz}\\geq\\rm10^{25}\\,W\\,Hz^{-1}$) and compact ($\\leq1$~kpc), characterized by a convex radio spectrum that peaks between 0.3-10~GHz (observer's frame), and they represent a significant fraction ($\\sim$ 10\\%) of the bright radio source population (O'Dea 1998). Only a few percent of them have weak extended emission (Stanghellini et al. 2005). A couple of GPS sources are also identified as Compact Symmetric Objects (CSOs, e.g Owsianik \\& Conway 1998). It is believed that GPS-small sizes are most likely due to their youth ($<{10^4}$ years) than to a dense confining medium (Murgia et al. 1999; O'Dea et al. 2005). In this scenario, these sources will evolve into large radio sources ($>15$~kpc), some of them would evolve into FRII radio sources (Fanti et al. 1995, Snellen et al. 2000). However, the GPS phenomenon is not completely understood, e.g. the distribution of sizes/ages, structure/spectra, and variability of GPS sources. We have carried out EVN (European VLBI Network) observations of 19 GPS sources, 15 of them are from the Parkes half-Jansky sample (Snellen et al. 2002) with declination $>-5^\\circ$ and not observed with VLBI before (except 2121$-$014 and 2322$-$040 we observed), four others are from our previous observation list which observed with the EVN at 2.3/8.4 GHz and/or 5 GHz (Xiang et al. 2005, 2006), the 1.6 GHz observation will further provide information on the source structure and spectra. Furthermore, we have carried out flux density observations with the Urumqi 25m telescope in order to study GPS flux density variability, searching for a different behavior between galaxies and quasars.   \\begin{table*}  \\caption[]{The GPS sources. Columns (1),(2) source names; (3) optical identification (gl: galaxy, Q: quasar, EF: empty field); (4) optical magnitude; (5) redshift (de Vries et al. 2007, those with * are photometric estimated by Tinti et al. 2005); (6) linear scale factor pc/mas [$H_{0}=71 km s^{-1} Mpc^{-1}$ and $q_{0}=0.5$ have been assumed]; (7) maximum angular size from the observation; (8) maximum linear size; (9) VLBI structure (cd: compact double, cj: core-jet, n: no detection); (10) low frequency spectral index; (11) higher frequency spectral index; (12) turnover frequency; (13) peak flux density; (14) references for the spectral information, 1 Snellen et al. 2002, 2 de Vries et al. 1997, 3 Stanghellini et al. 1998, where $S \\propto \\nu^{-\\alpha}$.} $$ \\begin{tabular}{cccccccccccccc}  \\hline \\noalign{\\smallskip} 1&2&3&4&5&6&7&8&9&10&11&12&13&14\\\\ $Name$ &other& $id$ & $m_{R}$ & $z$ & $pc/mas$ & $\\theta$ & $L$ & $vlbi$ & $\\alpha_{l}$ & $\\alpha_{h}$ & $\\nu_{m}$ & $S_{m}$ & $ref$ \\\\ & & & & && mas & pc &  & & & GHz & Jy & \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip}  J0210+0419 &B0208+040& gl   & $>{24.1}$  & 1.5* &6.1    & 90  &          & cd &  &0.80  & 0.4 & 1.3 & 1 \\\\  J0323+0534 &4C+05.14& gl & 19.2    & 0.1785 &2.7   &180 &     490 &  cd&  &0.85   & 0.4 & 7.1 & 1 \\\\  J0433$-$0229 & 4C$-$02.17&gl & 19.1  & 0.530 &  5.1  & 80 & 408          & cd  &  &0.52    & 0.4 & 3.0 & 1 \\\\  J0913+1454 & B0910+151&gl  & $22.9$ & 0.47* & 4.9   & 80  &            & cd &  &0.75  & 0.6 & 1.1& 1 \\\\  J1057+0012& B1054+004&gl   & $22.3$ & 0.65* & 5.5  & 80? &            & cj &     &0.67  & 0.4   & 1.6  & 1\\\\  J1109+1043& B1107+109&gl   & 22.6   & 0.55*  & 5.2   & 60  &            & cd  &  & 0.94  & 0.5 & 2.4& 1 \\\\  J1135$-$0021 &4C$-$00.45& gl  & 21.9  & 0.975 & 6.0 & 120 &  720          & cd  &   &0.78   & 0.4 &2.9 &1  \\\\  J1203+0414& B1200+045& Q   & 18.8  & 1.221 &  6.1     & 75 &   458      &  cj &  &0.45   & 0.4 &1.4 &1  \\\\  J1352+0232 &B1349+027& gl  & 20.0  &  0.607  & 5.4   & 170  & 918           &  cd&  &0.58   & 0.4 & 2.0 & 1 \\\\  J1352+1107 & 4C+11.46&gl  & $21.0$ & 0.891&  5.9 & 50 &   295         & cd  &  &1.03   & 0.4 & 3.6 & 1 \\\\  J1600$-$0037 & B1557$-$004 &gl &        &      &        &  50 &            & cd/cj    &  &1.17   & 1.0 & 1.2 & 1 \\\\  J1648+0242 &4C+02.43& gl   & $22.1$  & 0.824&  5.8  &    &           &  n    &   &   &0.4  &3.4 & 1 \\\\  J2058+0540 & 4C+05.78&gl   & 23.4  & 1.381 & 6.1   &160  &     970    & cd &   &0.95   & 0.4 &3.1  & 1 \\\\  J2123$-$0112 & B2121$-$014 & gl   & 23.3  & 1.158 & 6.1 & 80     &   488     & cso & -0.56 & 0.75  & 0.4 & 2.0 & 1 \\\\  J2325$-$0344 &B2322$-$040 & gl  & 23.5 &  1.509  &  6.0     &75&450 & cso  & -0.42 & 0.75  & 1.4 & 1.2 & 1 \\\\  J0917+1113 & B0914+114 & EF   &      &       &             & 190 &            & cso  & -0.1  & 1.6   & 0.3 & 2.3 & 3 \\\\  J1753+2750 & B1751+278 & gl  & 21.7  & 0.86* & 5.9        & 50   &              & cj  & -0.27 & 0.57  & 1.4 &0.6 &2 \\\\  J1826+2708 & B1824+271 & gl  & 22.9  &       &             & 45  &            &  cso & -0.39 & 0.75  & 1.0 & 0.4 & 2 \\\\  J2325+7917 & B2323+790 & gl   & 19.5V &       &             & 32  &            &  cd  &-0.3  &0.75   &1.4  &1.2 & 2 \\\\  \\noalign{\\smallskip} \\hline \\end{tabular}{} \\label{tb1} $$ \\end{table*} ",
        "conclusions": "Our VLBI results show that 80\\% of the GPS galaxies exhibit a compact double or CSO-like structure. However, for 13 sources which were only observed at 1.6 GHz further VLBI observations are needed for classifying their structure type in detail. Combining the 1.6, 2.3/8.4, and 5 GHz VLBI observations for other 6 sources, 4 CSOs are classified according to their steep spectrum of double lobes. The GPS galaxies are dominated by jet/lobe emission, with a hidden or weak core (it was not detected because of sensitivity in most of our sources), suggesting they are at large angles to the line of sight. While GPS quasars may be at moderate angles to the line of sight (Stanghellini et al. 2001).  Only 12\\% of GPS galaxies show flux variation higher than 10\\%, while 44\\% of GPS quasars show variability higher than 10\\% in passed 20 years. The result indicates that most GPS galaxies are stable in flux density, and the GPS quasars are much more variable than GPS galaxies. The variable GPS quasars often show compact core-jet, e.g. the quasar 0642+449 which has 140\\% flux variation in Fig.~\\ref{fig7} is an extremely compact core-dominated one at a redshift of 3.4. Torniainen et al. (2005) found that 54\\% of GPS sources show flux variability in long term (with fractional variability index $>3$). In our result, 25\\% of GPS sources show $>10\\%$ variation, the difference may be due to different samples and the definition of variability, we simply compared the flux densities in two epoches.  GPS galaxies which show comparable double lobes can be interpreted as type-2 AGN in the framework of unified scheme. GPS quasars, which show a core-jet structure and variability, could be mostly type-1 AGN."
    },
    "1005/1005.2685_arXiv.txt": {
        "abstract": "We apply our scalar-tensor-vector (STVG) modified gravity theory (MOG) to calculate the infall velocities of the two clusters constituting the Bullet Cluster 1E0657-06. In the absence of an applicable two-body solution to the MOG field equations, we adopt an approximate acceleration formula based on the spherically symmetric, static, vacuum solution of the theory in the presence of a point source. We find that this formula predicts an infall velocity of the two clusters that is consistent with estimates based on hydrodynamic simulations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.2961_arXiv.txt": {
        "abstract": "I report observations of unusually strong photospheric and chromospheric velocity oscillations in and near the leading sunspot of NOAA 10781 on 03 July 2005. I investigate an impinging wave as a possible origin of the velocity pattern, and the changes of the wave after the passage through the magnetic fields of the sunspot.  The wave pattern found consists of a wave with about 3 Mm {\\em apparent} wavelength that propagates towards the sunspot. This wave seems to trigger oscillations inside the sunspot's umbra, which originate from a location inside the penumbra on the side of the impinging wave. The wavelength decreases and the velocity amplitude increases by an order of magnitude in the chromospheric layers inside the sunspot. On the side of the sunspot opposite to the impinging plane wave, circular wave fronts centered on the umbra are seen propagating away from the sunspot outside its outer white-light boundary. They lead to a peculiar ring structure around the sunspot, which is visible in both velocity and intensity maps. The fact that only weak photospheric velocity oscillations are seen in the umbra -- contrary to the chromosphere where they peak there -- highlights the necessity to include the upper solar atmosphere in calculations of wave propagation through spatially {\\em and} vertically extended magnetic field concentrations like sunspots. ",
        "introduction": "Introduction} The effect of magnetic fields on the passage of (acoustic) waves is important for local helioseismology and time-distance analysis \\cite{cameron+etal2008,moradi+cally2008,khomenko+etal2009}. The magnetic fields of sunspots can change the characteristic mode of the wave and selectively absorb certain frequencies only \\cite{braun1995,cally2000}. Using the change of the wave pattern in and near sunspots, the properties of the solar velocity field even below the visible surface can be determined \\cite{duvall+etal1996,haber2008}. Whereas in most of these publications the effect of the magnetic fields on impinging waves was studied, some groups also have investigated the fate of waves that are generated inside a sunspot \\cite{khomenko+collados2006,olshevsky+etal2007,parchevsky+kosovichev2009}. \\inlinecite{khomenko+collados2006} found for instance that not only the wave mode changes during the propagation inside inclined magnetic fields, but also the propagation direction as well, up to a downward reflection of upwards propagating waves (see also \\opencite{khomenko+collados2009}). \\begin{figure} \\centerline{\\resizebox{9.8cm}{!}{\\hspace*{.2cm}\\includegraphics{fig1a.eps}}} \\centerline{\\resizebox{9.8cm}{!}{\\hspace*{.2cm}\\includegraphics{fig1b.eps}}} \\caption{Top row: LOS velocity of \\ion{Ti}{i} at 630.37\\thinspace nm in original resolution (left), smoothed by 15$\\times$15 pixels (middle), and smoothed and contrast enhanced (right). The black circle marks the outer limit of the penumbra. Lower row: continuum intensity at 630 nm, polarization degree, and LOS velocity of \\ion{He}{i} at 1083\\thinspace nm. For \\ion{He}{i} the displayed velocity range was doubled. Positive velocities correspond to blueshifts of the lines. The granulation contrast is about 3\\%; the difference between the intensity in the umbra ($\\approx$0.5 of I$_c$) and the quiet Sun ($\\approx$I$_c$) strongly suppresses its visibility in the image. The various white arrows denote the scan direction from left to right, the direction towards disc center, and towards solar North and East, respectively. Tick marks are in Mm.\\label{fig1}} \\end{figure}  The unusual velocity pattern near the leading sunspot of NOAA 10781 was actually detected only by chance. At the end of the usual reduction of data taken with the POlarimetric LIttrow Spectrograph \\cite{beck+etal2005b}, maps of the line-of-sight (LOS) velocity of the observed spectral lines are automatically created to be placed into an on-line overview archive (\\url{http://www3.kis.uni-freiburg.de/~cbeck/POLIS_archive/POLIS_archivemain.html}). The peculiarity of the observation in question went unnoticed up to that point, but a closer inspection of any velocity map of the POLIS lines from the data set in the archive or in the upper part of Figure 1 reveals at once the presence of bright and dark rings, most prominent in the upper left half of the map. The rings can be seen more clearly in the smoothed and contrast-enhanced velocity map (top middle and left). Especially if one demagnifies the image, as done automatically in the data archive, the visibility is strongly enhanced. The ring pattern does not seem to be related directly to photospheric magnetic fields; it is located partly at places without a polarization signal (middle bottom panel of Figure 1). The same pattern is visible also in the \\ion{Si}{i} line at 1082.7 nm observed simultaneously with the \\textit{ Tenerife Infrared Polarimeter} (TIP: \\opencite{martinezpillet+etal1999}; \\opencite{collados+etal2007}) while the chromospheric \\ion{He}{i} line at 1083 nm shows a markedly different pattern with strong oscillations in the umbra. Taken together, all spectral lines indicate unusually strong oscillations.  In the following section, I will discuss the observations in detail. I will then use different approaches to determine the spatial and temporal properties of the oscillations in the umbra and the surroundings of the sunspot in Section \\ref{ana}. I discuss a possible source of the observed  wave pattern in Section \\ref{disco}. The conclusions are presented in Section 5, while the Appendix A shows some numerical tests on the influence of the observation mode on the found wave properties. \\begin{table} \\caption{Spectral Lines Observed\\label{speclines} } \\begin{tabular}{cccc} \\multicolumn{2}{c}{POLIS} & \\multicolumn{2}{c}{TIP} \\cr \\ion{Fe}{i} & 630.15 nm & \\ion{Si}{i} & 1082.7 nm  \\cr \\ion{Fe}{i} & 630.25 nm & \\ion{He}{i} &  1083 nm\\cr \\ion{Ti}{i} & 630.37 nm &  & \\cr \\end{tabular} \\end{table} ",
        "conclusions": "The observed velocity pattern is unique in the sense that none of the several dozens sunspot observations with POLIS or TIP taken between 2003 and 2009 have shown a similar pattern with a comparable velocity amplitude. The wave pattern is directly visible in a single snapshot without the necessity of using difference images, opposite to the photospheric signature of the flare-induced ``sunquakes'' that generally show smaller velocity amplitudes. If the causal relationship developed in the previous section is  correct, then the strong magnetic fields of the sunspot have refracted a plane wave in the photosphere, which presumably was a common part of the solar acoustic $p$-mode oscillations, into a spherical acoustic wave, reducing the velocity amplitude but preserving the wavelength. As inside the umbra mainly the chromospheric helium line shows indications of strong oscillations, the wave should first have propagated to the upper atmosphere layers and then back to the photosphere inside the sunspot. \\inlinecite{khomenko+collados2006} found such an effect in their simulations, where a part of the wave energy was reflected downwards again, which fits with the reduction of the photospheric velocity amplitude seen in the observations here. This would imply that for modeling the wave propagation of oscillations inside the strong magnetic fields of sunspots also chromospheric atmosphere layers have to be included.  \\begin{acks} The VTT is operated by the Kiepenheuer-Institut f\\\"ur Sonnenphysik (KIS) at the Spanish Observatorio del Teide, operated by the Instituto de Astrof{\\'is}ica de Canarias (IAC). POLIS has been a joint development of the KIS and the High Altitude Observatory, Boulder, Colorado. The Dutch Open Telescope is operated by Utrecht University at the Spanish Observatorio del Roque de los Muchachos of the IAC. C.B.~acknowledges partial support by the Spanish Ministry of Science and Innovation through project AYA 2007-63881. I thank the PI of the ITP campaign in 2005, Y.~Katsukawa, for the permission to use the data. \\end{acks}  \\begin{appendix}"
    },
    "1005/1005.0964_arXiv.txt": {
        "abstract": "This paper presents the chemical abundance analysis of a sample of 18 giant stars in 3 old globular clusters in the Large Magellanic Cloud, namely NGC~1786, NGC~2210 and NGC~2257. The derived iron content is [Fe/H]=~--1.75$\\pm$0.01 dex ($\\sigma$=~0.02 dex), ~--1.65$\\pm$0.02 dex ($\\sigma$=~0.04 dex) and ~--1.95$\\pm$0.02 dex ($\\sigma$=~0.04 dex) for NGC~1786, NGC~2210 and NGC~2257, respectively. All the clusters exhibit similar abundance ratios, with enhanced values ($\\sim$+0.30 dex) of [$\\alpha$/Fe], consistent with the Galactic Halo stars, thus indicating that these clusters have formed from a gas enriched by Type II SNe. We also found evidence that {\\sl r}-process are the main channel of production of the measured neutron capture elements (Y, Ba, La, Nd, Ce and Eu). In particular the quite large enhancement of [Eu/Fe] ($\\sim$+0.70 dex) found in these old clusters clearly indicates a relevant efficiency of the {\\sl r}-process mechanism in the LMC environment. ",
        "introduction": "In the last decade, the advent of the high resolution spectrographs mounted on the 8-10 m telescopes has allowed to extend the study of the chemical composition of individual Red Giant Branch (RGB) stars outside our Galaxy up to dwarf and irregular galaxies of the Local Group. Chemical analysis of RGB stars are now available for several isolated dwarf spheroidal (dSph) galaxies as Sculptor, Fornax, Carina, Leo I, Draco, Sextans and Ursa Minor \\citep{shetrone01, shetrone03, letarte} and the Sagittarius (Sgr) remnant \\citep{boni00, monaco05, monaco07, sbordone}. As general clue, these studies reveal that the chemical abundance patterns in the extragalactic systems do not resemble those observed in the Galaxy, with relevant differences in the [$\\alpha$/Fe] \\footnote{We adopt the usual spectroscopic notations that [$X_1$/$X_2$]=~$\\lg(N_{X_{1}}/N_{X_{2}})_{*}$ - $\\lg(N_{X_{1}}/N_{X_{2}})_{\\odot}$ and that $\\lg(N_{X_{1}})$=~$\\lg(N_{X_{1}}/N_{H})$+12.}, [Ba/Fe] and [Ba/Y] ratios, thus suggesting different star formation history and chemical evolution \\citep[see e.g.][]{venn, geisler, tolstoy}.  Unlike the dSphs, the irregular galaxies as the Large Magellanic Cloud (LMC) contain large amount of gas and dust, showing an efficient ongoing star-formation activity. The LMC globular clusters (GCs) span a wide age/metallicity range, with both old, metal-poor and young, metal-rich objects, due to its quite complex star formation history. Several events of star formation occurred: the first one $\\sim$13 Gyr ago and 4 main bursts at later epochs, 2 Gyr, 500 Myr, 100 Myr and 12 Myr ago \\citep{harris09}. Until the advent of the new generation of spectrographs, the study of the chemical composition of the LMC stars was restricted to red and blue supergiants \\citep{hill95, korn00, korn02}, providing information only about the present-day chemical composition. The first studies based on high resolution spectra of RGB stars \\citep{hill00, johnson06, pompeia08} provided first and crucial  information about the early chemical enrichment and nucleosynthesis.   Of the $\\sim$300 compact stellar clusters listed by \\citet{kon}, metallicity determinations from Ca~II triplet are available for some tens of objects \\citep{ols91,aaron} and only for 7 clusters high-resolution spectroscopic analysis have been carried out \\citep{hill00, johnson06}. With the final aim of reconstructing the formation history of star clusters in the LMC, a few years ago we started a systematic spectroscopic screening of giants in a sample of LMC GCs with different ages.  In the first two papers of the series \\citep{f06,m08} we presented the chemical analysis of 20 elements for 4 intermediate-age LMC clusters (namely, NGC~1651, 1783, 1978, 2173). Moreover, \\citet{m09} discussed the iron content and the abundances of O, Na, Mg and Al for 3 old LMC clusters (namely NGC~1786, 2210 and 2257), discovering anticorrelation patterns similar to those observed in Galactic clusters. Here we extend the abundance analysis to additional 13 chemical elements in these 3 LMC clusters, also performing a detailed comparison with stellar populations in our Galaxy (both in the field and in globulars) and in nearby dSphs.  The paper is organized as follows: Section~2 presents the dataset and summarized the adopted procedure to derive radial velocities; Section~3  describes the methodology used to infer the chemical abundances; Section~4 discusses the uncertainties associated to the chemical abundances. Finally, Section~5 and ~6 present and discuss the results of the chemical analysis. ",
        "conclusions": "We have analyzed high-resolution spectra of 18 giants of 3 old LMC GCs, deriving abundance ratios for  13 elements, in addition to those already discussed in \\citet{m09}  and sampling the different elemental groups, i.e. iron-peak, $\\alpha$ and neutron-capture elements. The main results of our chemical analysis are summarized as follows: \\begin{itemize} \\item the three target clusters are metal-poor, with an iron content of [Fe/H]=~--1.75$\\pm$0.01 dex ($\\sigma$=~0.02 dex), ~--1.65$\\pm$0.02 dex ($\\sigma$=~0.04 dex) and ~--1.95$\\pm$0.02 dex ($\\sigma$=~0.04 dex) for NGC~1786, NGC~2210 and NGC~2257, respectively \\citep[see][]{m09}; \\item all the three clusters show the same level of enhancement of the $<\\alpha/Fe>$ ratio ($\\sim$~+0.30 dex), consistent with a gas enriched by type II SNe, while metal-rich, younger LMC clusters exhibit solar-scaled $<\\alpha/Fe>$ ratio, due to the contribution of type Ia SNe at later epochs; \\item the iron-peak elements (Sc, V, Cr, Ni)  follow a solar pattern (or slightly sub-solar, in some cases), according with the observed trend in our Galaxy and consistent with the canonical nucleosynthesis scenario; \\item the studied clusters show a relevant ($\\sim$--0.30 dex) depletion of [Y/Fe], while the other s-process elements (with the exception of Nd) display abundance ratios consistent with the Galactic distributions. [Ba/Fe] and [Ba/Y] in the old LMC GCs are lower than the values measured in the metal-rich, intermediate-age LMC GCs, because in the former the AGB stars had no time to evolve and enrich the interstellar medium; \\item ~[Eu/Fe] is enhanced ($\\sim$+0.70 dex) in all the clusters. This seems to suggest that the r-process elements production is very efficient in the LMC, being also the main channel of nucleosynthesis for the other neutron-capture elements. \\end{itemize}  In summary, the old, metal-poor stellar population of the LMC clusters closely resembles the GGCs in many chemical abundance patterns like the iron-peak, the $\\alpha$ and heavy {\\sl s-}process elements, and concerning the presence of chemical anomalies for Na, O, Mg and Al. When compared with dSphs the LMC old stellar population shows remarkably different abundance patterns for [$\\alpha$/Fe] and neutron-capture elements."
    },
    "1005/1005.2220_arXiv.txt": {
        "abstract": "{Active region moss are transition region phenomena, first noted in the images recorded by the Transition Region and Coronal Explorer (TRACE) in $\\lambda$171. Moss regions are thought to be the footpoints of hot loops (3-5~MK) seen in the core of active regions. These hot loops appear 'fuzzy' (unresolved). Therefore, it is difficult to study the physical plasma parameters in individual hot core loops and hence their heating mechanisms. Moss regions provide an excellent opportunity to study the physics of hot loops. In addition, they allow us to study the transition region dynamics in the footpoint regions.}{To derive the physical plasma parameters such as temperature, electron density, and filling factors in moss regions and to study their variation over a short (an hour) and a long time period (5 consecutive days).}{Primarily, we have analyzed spectroscopic observations recorded by the Extreme-ultraviolet Imaging Spectrometer (EIS) aboard Hinode. In addition we have used supplementary observations taken from TRACE and the X-Ray Telescope (XRT) aboard Hinode.}{The moss emission is strongest in the \\ion{Fe}{xii} and \\ion{Fe}{xiii} lines. Based on analyses using line ratios and emission measure we found that moss regions have a characteristic temperature of log~T[K]~=~ 6.2. The temperature structure in moss region remains almost identical from one region to another and it does not change with time. The electron densities measured at different locations in the moss regions using \\ion{Fe}{xii} ratios are about 1-3~$\\times$~10$^{10}$~cm$^{-3}$ and about 2-4~$\\times$~10$^{9}$~cm$^{-3}$ using \\ion{Fe}{xiii} and \\ion{Fe}{xiv}. The densities in the moss regions are similar in different places and show very little variation over short and long time scales. The derived electron density substantially increased (by a factor of about 3-4 or even more in some cases) when a background subtraction was performed. The filling factor of the moss plasma can vary between 0.1-1 and the path length along which the emission originates is from a few 100 to a few 1000 kms long. By combining the observations recorded by TRACE, EIS and XRT, we find that the moss regions correspond to the footpoints of both hot and warm loops.}{} ",
        "introduction": " ",
        "conclusions": "\\label{con}  Using Hinode/EIS observations, we have studied basic physical plasma parameters such as temperature, electron density, filling factors, and column depth in moss regions and the variation of these parameters over an hour and over a time period of five days. In addition, we have revisited the question of whether the moss regions are the footpoints of hot loops using observations from TRACE, EIS and XRT. The results are summarized below.  \\begin{itemize}  \\item Based on the TRACE, EIS and XRT observations we find that most of the moss regions are essentially located at the footpoints of hot loops. In some places we observed TRACE~171 (1~MK) loops rooted in the moss regions.  \\item Based on the line intensity ratios of \\ion{Fe}{xi}~$\\lambda$188 and \\ion{Fe}{xiii}~$\\lambda$202, and an emission measure analysis, we find that the characteristic temperature of moss regions is about log~T~=~6.2. Emission measure analyses over a time period of one hour (Fig.~\\ref{emission_measure}) and over five days (Fig.~\\ref{five_days_em}) reveal that the thermal structure of the moss regions does not change significantly with time.  \\item The electron densities measured using \\ion{Fe}{xii} ratios are about 1-3~$\\times$~10$^{10}$~cm$^{-3}$ and about 2-4~$\\times$~10$^{9}$~cm$^{-3}$ using \\ion{Fe}{xiii} and \\ion{Fe}{xiv}. Work is in progress to try to resolve this discrepancy. It is worth emphasizing here that if the \\ion{Fe}{xii}~$\\lambda$186.8 were blended and we lower its intensity by 20\\%, then the electron densities obtained using \\ion{Fe}{xii}  would become consistent with those obtained using \\ion{Fe}{xiii} and \\ion{Fe}{xiv}. The densities derived using \\ion{Fe}{xiii} and \\ion{Fe}{xiv} are similar to those derived by \\cite{fletcher_moss} using \\ion{Si}{x} line ratios observed by CDS. However, when we subtract the foreground/background emission we find a substantial increase (a factor of 3-4 or even more in some cases) in the densities.  \\item The electron densities only show small changes ($\\sim$25\\%) over a period of an hour. There are large variations (an order of magnitude increase) in \\ion{Fe}{xii} densities when measured over a period of five days. However, the variation in the densities obtained using \\ion{Fe}{xiii} and \\ion{Fe}{xiv} is only about 50-70\\%.  \\item The filling factor of the moss plasma is in the range 0.1-1 and the path length along which the emission originates is from a few 100 to a few 1000 kms long.  \\end{itemize}  These new measurements of the thermal and density structure in moss regions should provide important constraints for the modelling of loops in the hot and dense core of active regions."
    },
    "1005/1005.5636_arXiv.txt": {
        "abstract": "The Hubble tuning fork diagram, based on morphology and established in the 1930s, has always been the preferred scheme for classification of galaxies. However, the current large amount of data up to higher and higher redshifts asks for more sophisticated statistical approaches like multivariate analyses. Clustering analyses are still very confidential, and do not take into account the unavoidable characteristics in our Universe: evolution. Assuming branching evolution of galaxies as a 'transmission with modification', we have shown that the concepts and tools of phylogenetic systematics (cladistics) can be heuristically transposed to the case of galaxies. This approach that we call ``astrocladistics'', has now successfully been applied on several samples of galaxies and globular clusters. Maximum parsimony and distance-based approaches are the most popular methods to produce phylogenetic trees and, like most other studies, we had to discretize our variables. However, since astrophysical data are intrinsically continuous, we are contributing to the growing need for applying phylogenetic methods to continuous characters.  classification; astrophysics; cladistics; multivariate; evolution; galaxies; continuous characters ",
        "introduction": "The extragalactic nature of galaxies was discovered by Edwin Hubble less than 90 years ago. He also discovered the expansion of the Universe and established in 1936 the famous tuning fork diagram or Hubble diagram based on the idea that elliptical galaxies should evolve into flattened systems such as ordinary or barred spiral galaxies. Since then, telescopes and the associated detectors have made remarkable progress. We are now able to study galaxies in great detail, identifying individual stars, gas and dust clouds, as well as different stellar populations. Imagery brings very fine structural details, and spectroscopy provides the kinematical, physical and chemical conditions of the observed entities at different locations within the galaxy. For more distant objects, information is scarcer, but deep systematic sky surveys gather millions of spectra for millions of galaxies at various redshifts.  Like paleontologists, we observe objects from the distant past (galaxies at high redshift), and like evolutionary biologists, we want to understand their relationships with nearby galaxies, like our own Milky Way. Strangely enough, the Hubble classification, depicted in the Hubble diagram, is still frequently used as a support to describe galaxy evolution, even though it ignores all observables except morphology (Fraix-Burnet et al 2006a, Hernandez and Cervantes-Sodi 2006). Can the enormous amount of very detailed observations and galaxy diversity be solely depicted by few large families characterized only by their global shape? There obviously must be a better way to exploit the data gathered by very large telescopes and their sophisticated detectors, and to account for the complicated physical processes that lead to the diversification of galaxy.  Abandoning the one-parameter classification approach and using all available descriptors means taking a methodological step equivalent to the one biologists took after Adanson and Jussieu in the 18th century. One basic tool, the Principal Component Analysis, is relatively well-known in astrophysics (e.g. Cabanac et al 2002, Recio-Blanco 2006). However, only a very few attempts to apply multivariate clustering methods have been made very recently (Chattopadhyay and Chattopadhyay 2006, 2007; Chattopadhyay et al 2007, 2008, 2009). Sophisticated statistical tools are used in some areas of astrophysics and are developing steadily, but multivariate analysis and clustering techniques have not much penetrated the community.  A supplementary difficulty is that evolution, an unavoidable fact, is not correctly taken into account in most classification methods. By mixing together objects at different stages of evolution, most of the physical significance and usefulness of the classification is lost. In practice, the evolution of galaxies is often limited to the evolution of the properties of the entire population as a function of redshift (Bell 2005). Since environment (the expanding Universe) and galaxy properties are so much intricate, this kind of study is relevant to a first approximation. However, recent observations have revealed that galaxies of all kinds do not evolve perfectly in parallel, as illustrated for instance by the so-called downsizing effect which shows that large galaxies formed their stars earlier than small ones (e.g. Neistein 2006). New observational instruments now brings multivariate information at different stages of evolution, and in various evolutive environments. In this multivariate context, we believe that the notion of \"evolution\", easy to understand for a single parameter, is advantageously replaced by \"diversification\".  In this conference, I present our efforts to implement a method to reconstruct the \"galactogenesis\", based on tools and concepts largely developed in evolutionary biology and bioinformatics. I first present cladistics and some of our results. Then I replace this approach among multivariate statistical methods and explain why we also need these other tools to explore the parameter space. Finally, I present one of our developments that makes the astrophysical problematics to join current bioinformatics and mathematics studies about the use of continuous characters in reconstructing phylogenies. ",
        "conclusions": "It is now clear that cladistics can be applied and be useful to the study of galaxy diversification. Many difficulties, conceptual and practical, have been solved,. Significant astrophysical results have been obtained and will be extended to larger samples of galaxies and globular clusters. However, many paths remain in the exploration of this new and large field of research.  There are difficulties that seem to be intrinsic to astrophysics. Most notably, we have millions of objects, but a few tens of descriptors. Of course the situation will improve with time, in particular with integral-field spectroscopy (spectra of detailed regions in the galaxy, e.g. Ensellem et al. 2007). Spectra might not be currently employed at their full capacity of description. Here also, sophisticated statistics can help. But it is not clear whether these data would lead to the discrimination of hundreds of classes. Perhaps this is an erroneous target, perhaps galaxies cannot be classified with such a refinement. Nonetheless, this is already a matter of using multivariate clustering methods and interpreting their results usefully. This is probably more than a conceptual question because clustering with continuous data and large intraspecific variance is a very complicated problem in itself. Sophisticated statistical tools must be used, but the question of characterizing the groups in this context requires a different culture that is not yet widespread in astrophysics. We are convinced that improvements can be made here.  Cladistics is supposed to identify clades, that are evolutionary groups, whereas the concept of \u201cspecies\u201d is not defined at all in astrophysics, and we have even not converged toward groupings based on multivariate analyses. Understanding the formation and evolution of galaxies, depicting a global scenario for galaxy diversification, is a major challenge for the astrophysics of XXIst century."
    },
    "1005/1005.2827_arXiv.txt": {
        "abstract": "Mass accretion onto (proto-)stars at high accretion rates $\\dot{M}_\\ast > 10^{-4}~M_{\\odot}~{\\rm yr}^{-1}$ is expected in massive star formation. We study the evolution of massive protostars at such high rates by numerically solving the stellar structure equations. In this paper we examine the evolution via disk accretion. We consider a limiting case of ``cold'' disk accretion, whereby most of the stellar photosphere can radiate freely with negligible backwarming from the accretion flow, and the accreting material settles onto the star with the same specific entropy as the photosphere. We compare our results to the calculated evolution via spherically symmetric accretion, the opposite limit, whereby the material accreting onto the star contains the entropy produced in the accretion shock front. We examine how different accretion geometries affect the evolution of massive protostars. For cold disk accretion at $10^{-3}~M_{\\odot}~{\\rm yr}^{-1}$ the radius of a protostar is initially small, $R_\\ast \\simeq$ a few $R_\\odot$. After several solar masses have accreted, the protostar begins to bloat up and for $M_\\ast \\simeq 10~M_\\odot$ the stellar radius attains its maximum of $30 - 400~R_\\odot$. The large radius $\\sim$$100~R_\\odot$ is also a feature of spherically symmetric accretion at the same accreted mass and accretion rate. Hence, expansion to a large radius is a robust feature of accreting massive protostars. At later times the protostar eventually begins to contract and reaches the Zero-Age Main-Sequence (ZAMS) for $M_\\ast \\simeq 30~M_\\odot$, independent of the accretion geometry. For accretion rates exceeding several $10^{-3}~M_{\\odot}~{\\rm yr}^{-1}$ the protostar never contracts to the ZAMS. The very large radius of several $100{\\rm s}~R_\\odot$ results in a low effective temperature and low UV luminosity of the protostar. Such bloated protostars could well explain the existence of bright high-mass protostellar objects, which lack detectable H~II regions. ",
        "introduction": "\\label{sec:intro}  Recent studies have revealed that the formation process of massive $(> 8~M_\\odot)$ stars differ in many respects from that of low-mass $(\\sim 1~M_\\odot)$ stars (e.g., Zinnecker \\& Yorke 2007). Although signatures of mass accretion are widely observed toward forming massive stars, the estimated accretion rates usually exceed $10^{-4}~M_{\\odot}~{\\rm yr}^{-1}$, which are much higher than the typical value in low-mass star formation, \\begin{equation} \\dot{M}_\\ast \\sim \\frac{c_s^3}{G} \\simeq 5 \\times 10^{-6} \\left( \\frac{T}{10~{\\rm K}} \\right)~ M_\\odot~{\\rm yr}^{-1}, \\label{eq:lmass_acc} \\end{equation} where $c_s$ and $T$ is the sound speed and temperature in natal molecular cloud cores, and $G$ is the gravity constant. In the past it has been argued that such high accretion rates are required for accretion flow to overcome the strong radiation pressure from forming massive stars (e.g., Larson \\& Starrfield 1971; Wolfire \\& Cassinelli 1987). However, these results were based on the assumption of spherically symmetric infall.  Nevertheless, there are other good reasons to expect high accretion rates during the formation of massive stars: Massive stars have short Kelvin-Helmholz time scales, once formed they consume the available nuclear fuel quickly, and they will lose a significant part of their mass through strong stellar winds.  Thus, massive stars evolve quickly, even while accreting.  We can expect the time spent in the main accretion phase to be significantly less than the main sequence lifetime, presumably of the order of a few dynamical times of the star-forming molecular core.  For a $30~M_\\odot$ star this provides a lower limit to the average accretion rate $\\simeq 10^{-4}~M_{\\odot}~{\\rm yr}^{-1}$. Accretion rates may vary strongly, however. Thus, phases of even higher accretion rates are likely.  The high accretion rates suggest that the nature of massive star formation differs from that of low-mass star formation. So far, various formation scenarios, such as monolithic collapse of massive dense cores (e.g., Yorke \\& Sonnhalter 2002; McKee \\& Tan 2002, 2003; Krumholz et al. 2007, 2009), and the competitive accretion with global infall of the cluster-forming clumps (e.g., Bonnell et al. 1998; Bonnell, Vine \\& Bate 2004; Wang et al. 2010), have been proposed. The proposed formation scenarios are being assessed with high-resolution observations of the massive cores and clumps (e.g., Motte et al. 2007; Marseille et al. 2008; Zhang et al. 2009; Bentemps et al. 2010).  The evolution of massive protostars at such high rates is a critical feature of massive star formation which determines the effect these stars have on their environment (e.g., Zinnecker \\& Yorke 2007). In our previous work (Hosokawa \\& Omukai 2008; 2009, hereafter Paper I), we examined the evolution by numerically solving the interior structure of a protostar with a dusty accretion envelope under the assumption of spherical symmetry. Our calculations show that accreting massive protostars have certain characteristic features. At $\\dot{M}_\\ast = 10^{-3}~M_{\\odot}~{\\rm yr}^{-1}$, for example, the protostellar radius becomes very large, exceeding $100~R_\\odot$ at maximum. The evolution before the arrival to the zero-age main sequence (ZAMS) lasts until $M_\\ast \\simeq 30~M_\\odot$. At even higher accretion rates $\\dot{M}_\\ast > 3 \\times 10^{-3}~M_{\\odot}~{\\rm yr}^{-1}$, the protostar very abruptly inflates and overtakes material in the inner part the accreting envelope, thus invalidating the assumption of steady mass accretion before reaching the ZAMS.  The very large stellar radius leads to a low effective temperature of the protostar, which in turn explains some observational properties of massive protostars in Orion KL nebula (Morino et al. 1998; Furuya \\& Shinnaga 2009).  Observations suggest that mass accretion onto massive stars is not spherically symmetric, but rather proceeds through massive circumstellar disks (e.g., Cesaroni et al. 2007). Signatures of rotating infall are widely found toward forming massive stars, sometimes associated with massive outflows along the rotation axis (e.g., Patel et al. 2005; Beltr\\'an et al. 2006a). Recent numerical simulations also suggest that the massive disk is a natural outcome of collapse of a massive dense core (Yorke \\& Sonnhalter 2002; Krumholz et al. 2007, 2009). Protostellar evolution via disk accretion has been studied for lower-mass protostars with low accretion rates $\\dot{M}_\\ast \\leq 10^{-4}~M_{\\odot}~{\\rm yr}^{-1}$ (e.g., Palla \\& Stahler 1992; Beech \\& Mitalas 1994; Hartman et al. 1997). These authors considered the limiting case of ``cold'' disk accretion, whereby most of the stellar surface is able to radiate into free space and the accreting material brings only a small amount of entropy into the star. They showed that, with the low accretion rates, disk accretion slightly reduces the stellar radius compared to the spherically symmetric cases.  Maeder and coworkers studied the evolution of accreting massive protostars in the cold disk limit with mass accretion rates increasing with the stellar mass up to $10^{-4}~M_\\odot~{\\rm yr}^{-1}$ (e.g., Norberg \\& Maeder 2000; Behrend \\& Maeder 2001). Recently, Yorke \\& Bodenheimer (2008, hereafter YB08) calculated the evolution in the same limit with much higher accretion rates $10^{-3}$ and $10^{-2}~M_\\odot~{\\rm yr}^{-1}$. Their calculated evolution differs from the spherical accretion case in an early phase of $M_\\ast \\lesssim 8~M_\\odot$. The stellar radius is several 10s$~R_\\odot$ with spherical accretion, whereas it is a few $R_\\odot$ with disk accretion. However, YB08 find that the stellar radius rapidly increases soon afterwards and exceeds $100~R_\\odot$ at $M_\\ast \\simeq 10~M_\\odot$. The subsequent evolution at $M_\\ast \\gtrsim 10~R_\\odot$ is quite similar for both spherical accretion and disk accretion.  In this paper, we examine how the different accretion geometries affect the evolution of massive protostars at high accretion rates. We explain why the two extreme cases lead to the same characteristic features of massive protostars, such as the large radius, in spite of the different evolution in the early phases. To this end, we calculate the protostellar evolution in the cold disk accretion limit with the numerical code used in Paper I. This allows us to extract how the protostellar evolution depends on the different accretion geometry independent of the effects of different stellar evolution codes.  For completeness, we also compare the calculated evolution with previous results by YB08.  Finally, we stress that this study is very relevant to star formation in the early universe, where high accretion rates are also expected. Protostellar evolution with high accretion rates has been studied with zero or very low metallicities under the assumption of spherically symmetric accretion (e.g., Stahler, Palla \\& Salpeter 1986; Omukai \\& Palla 2001, 2003; Hosokawa \\& Omukai 2009b). Recent numerical simulations show that circumstellar accretion disks are also expected during star formation in the early universe (Yoshida, Omukai \\& Hernquist 2008; Stacy, Grief \\& Bromm 2010). Tan \\& McKee (2004) examined the evolution of primordial protostars via disk accretion using simple analytic models. Here, we calculate the detailed evolution of a primordial protostar in the cold disk accretion limit.  The organization of this paper is as follows. The numerical method and the cases considered are briefly described in Section \\ref{sec:num}. In Section \\ref{sec:result}, we focus on the protostellar evolution at the accretion rate $\\dot{M}_\\ast = 10^{-3}~M_{\\odot}~{\\rm yr}^{-1}$, which allows a direct comparison to the results of YB08. First, we briefly review the evolution via spherical accretion studied in Paper I (Section \\ref{ssec:sp_acc}), and then consider the evolution via cold disk accretion (Section \\ref{ssec:d_acc}). Evolution of a primordial protostar is investigated in Section \\ref{ssec:metal} and subsequently contrasted to the present-day cases. Variation of the protostellar evolution with different accretion rates is studied in Section \\ref{sec:mdot_dep}. Finally, Section \\ref{sec:sum} is devoted to our summary and conclusions.  \\begin{figure*} \\plotone{f1.eps} \\caption{Schematic figures of a protostar with different accretion geometry: (a) spherically symmetric accretion, and (b) cold disk accretion. In the spherical case (a), the accretion flow directly hits the stellar surface forming an accretion shock front. If the flow is optically thick before reaching the surface, the photosphere locates outside the stellar surface. Dust grains in the accretion envelope evaporate in a dust destruction front far outside the photosphere. In the cold disk accretion case (b), gas predominantly accretes onto a circumstellar disk rather than the star. Accreting material settles onto the stellar surface through a geometrically thin layer (or possibly through geometrically thin accretion columns -- not shown).  Heat brought into the star in the accretion flow radiates freely into space until the temperature attains the photospheric value. Most of the stellar surface is unaffected by the accretion flow. The energy radiated away by the disk and/or accretion columns before the material settles onto the star (the so-called ``accretion luminosity'') needs to be accounted for separately from the intrinsic stellar luminosity.} \\label{fig:prost_schem_spph} \\end{figure*} ",
        "conclusions": "\\label{sec:sum}  In this paper we have studied the evolution of massive protostars with disk accretion at high rates $\\dot{M}_\\ast > 10^{-4}~M_{\\odot}~{\\rm yr}^{-1}$. We considered the limiting case of ``cold'' disk accretion, whereby most of the stellar surface is not affected by the accretion flow, and the accreting material brings a minimum of entropy into the star. We calculated the evolution of protostars in this limiting case by adopting the photospheric (freely radiating) boundary condition. The calculated evolution was compared and contrasted to the evolution with spherically symmetric accretion, corresponding to the opposite limit, whereby the accreting gas transports the maximum amount of entropy into the star.  First, we considered in detail the evolution for $\\dot{M}_\\ast = 10^{-3}~M_{\\odot}~{\\rm yr}^{-1}$. The basic evolution via the cold disk accretion is summarized as follows:  The entire evolution is divided into four evolutionary phases: (I) convection, (II) swelling, (III) KH contraction, and (IV) the main-sequence accretion phase. The evolution in the first three phases varies with different initial models. Roughly speaking, the stellar radius reaches its maximum $30 - 400~R_\\odot$ for $M_\\ast \\simeq 10~M_\\odot$ at the end of the swelling phase (II). On the other hand, the protostar begins the final phase (IV) for $M_\\ast \\simeq 30~M_\\odot$, independent of different initial models.  A key physical quantity governing the evolution is the opacity in the stellar interior, analogous to the spherical accretion case. The opacity decreases with increasing stellar mass, which causes the variety of the evolution. In phase (I) the opacity is so high that the radiative heat transport within the star does not affect the stellar structure. In phase (II) the entropy accumulated in the stellar interior is gradually transported outward as the opacity decreases. A part of the transported entropy is received in a thin layer near the stellar surface. Due to the high entropy deposited into this layer it significantly bloats up, resulting in the swelling seen in phase (II). Afterwards, the opacity decreases sufficiently, allowing the star to lose heat by efficiently radiation. The protostar contracts and the interior densities and temperatures rise. This is the KH contraction phase (III). The protostar finally reaches the main-sequence accretion phase (IV), just after the active hydrogen burning begins.  A major difference between the results for spherical accretion and cold disk accretion is the role of deuterium burning. In the limit of cold disk accretion the accreting gas brings a minimal amount of entropy into the star. Since the stellar radius is smaller because of the lower average entropy in the stellar interior, the protostellar radius is initially small, $R_\\ast \\simeq$ a few $\\times R_\\odot$. However, deuterium burning begins earlier, which raises the entropy content within the star. This enhances the swelling. Consequently, the maximum radius is as large as that with spherical accretion.  Next, we calculated the protostellar evolution at zero metallicity in the cold disk accretion limit. The protostellar evolution varies with metallicities in the same manner as with spherical accretion; at $Z = 0$, the protostar swells up slightly earlier, and reaches the ZAMS with a smaller stellar radius.  Finally, we examined the evolution with different accretion rates. Our calculations show some dependences on accretion rates which were also found with spherical accretion. First, the stellar mass at the arrival to the ZAMS is higher for higher accretion rates. However, if the accretion rate is higher than a few $10^{-3}~M_\\odot~{\\rm yr}^{-1}$, the stellar radius continues to increase without returning to KH contraction. The protostar never reaches the ZAMS by steady mass accretion at such high rates.  The fact that a massive protostar bloats up to a radius of several 100s$~R_\\odot$ with an accretion rate exceeding $10^{-3}~M_\\odot~{\\rm yr}^{-1}$ is a very robust result of our studies. The large radius leads to a low effective temperature and a correspondingly low stellar UV luminosity. Thus, the growth of an H~II region will be delayed until mass accretion onto the star deceases significantly. This explains the existence of HMPOs with high luminosities exceeding $10^4~L_\\odot$.   {"
    },
    "1005/1005.0007_arXiv.txt": {
        "abstract": "{} {This work studies in more detail the stellar population, including its photometric properties and characteristics, in the rarely studied southern Galactic globular cluster NGC 1261. We focus on the brighter sequences of the cluster's color-magnitude diagram (CMD). Like in our previous works, we rely upon photometry in several passbands to achieve more reliable results and conclusions.} {We carried out and analyzed new multi-color photometry of NGC 1261 in $UBVI$ reaching below the turnoff point in all passbands in a fairly extended cluster field, about 14$\\arcmin$x14$\\arcmin$.} {We found several signs of the inhomogeneity (\"multiplicity\") in the stellar population. The most prominent of them are: (1) the dependence of the radial distribution of sub-giant branch (SGB) stars in the cluster on their $U$ magnitude, with brighter stars less centrally concentrated at the 99.9 \\% level than their fainter counterparts; (2) the dependence of the location of red giant branch (RGB) stars in the $U$-$(U-B)$ CMD on their radial distance from the cluster center, with the portion of stars bluer in the $(U-B)$ color increasing towards the cluster outskirts. Additionally, the radial variation of the RGB luminosity function in the bump region is suspected. We assume that both the SGB stars brighter in the $U$ and the RGB stars bluer in the $(U-B)$ color are probably associated with blue horizontal branch stars, because of a similarity in their radial distribution in the cluster. We estimated the metalicity of NGC 1261 from the slope of the RGB in $U$-based CMDs and the location of the RGB bump on the branch. These metallicity indicators give [Fe/H]$_{ZW} = -1.34 \\pm 0.16$ dex and [Fe/H]$_{ZW} = -1.41 \\pm 0.10$ dex, respectively. We isolated 18 probable blue straggler candidates. They are more centrally concentrated than the lower red giants of comparable brightness at the 97.9 \\% level. Their photometric characteristics imply that their majority is not consistent with the collisional origin. We also reliably isolated the asymptotic giant branch (AGB) clump and estimated the parameter $\\Delta V_{ZAHB}^{clump}$ = 1.01 $\\pm$ 0.06, that is the difference between the $V$-levels of the zero age HB and the clump.} {}      \\keywords {globular clusters: general -- globular clusters: individual: NGC 1261} ",
        "introduction": "\\label{introduc}  The southern Galactic globular cluster (GC) NGC 1261 ($\\alpha_{2000} = 03^{h}12^{m}15\\fs3$ and $\\delta_{2000} = -55\\degr13\\arcmin01\\arcsec$) remains one of the rarely studied GCs up to now, in spite of its moderate central concentration ($c$ = 1.27), negligible or no reddening because of the relatively high Galactic latitude ($b$ = $-52\\fdg13$), and the somewhat brighter absolute visual magnitude (M$_{V}$ = -7.81) compared to the most probable value of this characteristic of Galactic GCs (Harris \\cite{harris96}). Although the history of CMD studies of this object spans almost 40 years, it resulted so far in only six publications. The first, scattered photographic CMD reaching the horizontal branch (HB) of the cluster, was obtained by Alca\\'ino \\& Contreras (\\cite{alcontrer71}) and another one of a higher quality, by Alca\\'ino (\\cite{alcaino79}). The first CCD photometry of NGC 1261, aimed at estimating the cluster's age, was carried out twenty years ago by Bolte \\& Marleau (\\cite{boltemar89}). Since then three photometric studies were published in 1990s: Alca\\'ino et al. (\\cite{alcainoetal92}) focused mainly on the age determination, based on photometry in three bands; Ferraro et al. (\\cite{ferraroetal93}) studied the brighter sequences of the CMD, and Zoccali et al. (\\cite{zoccalietal}) obtained the luminosity and mass functions of NGC 1261. Yet, Hubble Space Telescope photometry in the cluster has been carried out by Piotto et al. (\\cite{piottoetal}), among 74 Galactic GCs.  Note than our short overview is not concerned with publications dealing with variable stars in NGC 1261 and other particular issues.  In this work, we study for the first time the stellar population of NGC 1261 based on photometry that covers a wide cluster field (approximately $14\\arcmin \\times 14\\arcmin$) and was conducted in several bands, including measurements in the $U$ bandpass, which is rarely used in investigations of Galactic GCs. In particular, we estimated the cluster metalicity by applying metalicity indicators not previously used in NGC 1261, one of which is related to the $U$-based CMDs. Also, signs of a probable inhomogeneity of the stellar population in NGC 1261 were found with $U$ photometry in the cluster.   \\begin{figure*} \\centerline{ \\includegraphics[clip=,angle=0,width=8 cm]{13449fg1a.eps} \\includegraphics[clip=,angle=0,width=8 cm]{13449fg1b.eps} } \\caption{Comparison of our present photometry of NGC 1261 and the photometry by Stetson (\\cite{ste00}) of standard stars in the cluster, taken from the Stetson Photometric Standard Fields; the differences in magnitudes and colors are derived by our data minus those of Stetson.} \\label{comparphot} \\end{figure*} ",
        "conclusions": "We obtained new $UBVI$ photometry of 5\\,481 stars in a wide field, approximately 14$\\arcmin$x14$\\arcmin$, of the southern, rarely studied GC NGC 1261 below its TO point in all passbands. We exploited this multi-color photometry especially in $U$ passband and its advantages over two-color photometry to further study this GC.  We found two very probable indications and one suspected sign of the inhomogeneity (multiplicity) of the stellar population in NGC 1261. First, there is an obvious dependence of the radial distribution of SGB stars in the cluster on their $U$ magnitude: brighter stars are less centrally concentrated at the 99.9 \\% confidence level and less numerous than their fainter counterparts. Second, RGB stars exhibit a systematically different location in the $U$-$(U-B)$ diagram at different radial distance from the cluster center: the proportion of stars bluer in the $(U-B)$ (or/and brighter in the $U$) increases towards the cluster outskirts. Additionally, the distribution of stars on magnitude, particularly in the $U$ band, in the RGB bump region is suspected to change radially in the cluster. We conclude that both the SGB stars brighter in $U$ and the RGB stars bluer in the $(U-B)$ may be associated with the blue BHB, because they are less numerous and less concentrated to the cluster center as compared to their counterparts in the respective evolutionary sequences.  We estimated the metalicity of NGC 1261 by applying metalicity indicators, which were previously not used on this cluster. One of them was recently proposed by us. It is related to deep, $U$-based CMDs of GCs that simultaneously reach the TO point and trace the upper RGB well with a sufficient amount of stars. The given condition was accomplished by our photometry. We obtained [Fe/H]$_{ZW} = -1.34 \\pm 0.16$ dex. We also found [Fe/H]$_{ZW} = -1.41 \\pm 0.10$ by exploiting the dependence of the position of the RGB bump on metalicity. Both estimates agree very well with each other. The resulting mean value of the cluster's metalicity, [Fe/H]$_{ZW} = -1.38 \\pm 0.14$ dex, is at the metal-poor extreme of the range of previously made estimates on the same metalicity scale obtained by different methods and indicators.  We isolated 18 BS candidates, a quite small population of these stars. They are found to be more centrally concentrated than the lower red giants of a comparable brightness. This difference is statistically significant at the 97.9 \\% confidence level according to the K-S test. Their position in the $U$-based two-color diagram indicates that the formation of their majority is not due to the collisional mechanism."
    },
    "1005/1005.1658.txt": {
        "abstract": "We have used multi-band imaging to investigate the nature of an extreme starburst environment in the nearby Lyman-break galaxy analogue Haro 11 (ESO\\,350-IG038) by means of its stellar cluster population. The central starburst region has been observed in 8 different high resolution {\\it HST} wavebands, sampling the stellar and gas components from UV to NIR. Photometric imaging of the galaxy was also carried out at 2.16 $\\mu$m by NaCo AO instrument at the ESO Very Large Telescope (VLT). We constructed integrated spectral energy distributions (SEDs) for about 200 star clusters located in the active star-forming regions and compared them with single stellar population models (suitable for physical properties of very young cluster population) in order to derive ages, masses and extinctions of the star clusters. The cluster age distribution we recover confirms that the present starburst has lasted for 40 Myr, and shows a peak of cluster formation only 3.5 Myr old. With such an extremely young cluster population, Haro 11 represents a unique opportunity to investigate the youngest phase of the cluster formation process and evolution in starburst systems. We looked for possible relations between cluster ages, extinctions, and masses. Extinction tends to diminish as function of the cluster age, but the spread is large and reaches the highest dispersion for clusters in partial embedded phases ($<5$ Myr). A fraction of low-mass (below $10^4 \\msun$), very young ($1-3$ Myr) clusters is missing, either because they are embedded in the parental molecular cloud and heavily extinguished, or because of blending with neighbouring clusters. The range of the cluster masses is wide, we observe that more than 30 \\% of the clusters have masses above $10^5 \\msun$, qualifying them as super star clusters. Almost half of the cluster sample is affected by flux excesses at wavelengths $> 8000$ \\AA\\ which cannot be explained by simple stellar evolutionary models. Fitting SED models over all wavebands leads to systematic overestimates of cluster ages and incorrect masses for the stellar population supplying the light in these clusters. We show that the red excess affects also the {\\it HST} F814W filter, which is typically used to constrain cluster physical properties. The clusters which show the red excess are younger than 40 Myr; we discuss possible physical explanations for the phenomenon. Finally, we estimate that Haro 11 has produced bound clusters at a rate almost a factor of 10 higher than the massive and regular spirals, like the Milky Way. The present cluster formation efficiency is $\\sim 38$\\% of the galactic star formation rate. ",
        "introduction": "Young star clusters (SCs) are luminous and quite compact objects which in some ways resemble older massive globular clusters. However while these old systems are \"passively\" fading due to stellar and secular evolution, young SCs are still, or in most of the cases, embedded in their dusty cocoons and at constant risk of destruction by interactions with the interstellar medium, internal feedback, tidal forces, and evaporation (see \\citealp{2009arXiv0911.0796L} and references therein). The youngest phases, when the cluster is still partially embedded in its natal molecular cloud, are difficult to explore because of the high extinction levels which introduce uncertainties in age and mass estimates. \\citet{2003ARA&A..41...57L} observed that the embedded phase in young Milky Way (MW) clusters lasts a few million years. Very young clusters in the Antennae (\\citealp{2002AJ....124.1418W}) system, with ages $\\leq$ 4 Myr, show a wide range of extinction between $0.5 \\leq A_V \\leq 7.6$ mag. \\citealp{2002AJ....124.1418W} identified an additional population of \"red clusters\", which they interpreted as young embedded clusters. Among them, the 2 Myr old WS\\,80, with its extreme $V-I = 2.92$ colour, mass $4 \\times 10^6 \\msun$ and $A_V = 7.6$ mag, is the brightest infrared source in the Antennae system. \\citet{2002AJ....124.1418W} indicated roughly 6 Myr as the age at which a young cluster has gotten rid of the dusty material left over from its parent cloud. In another galaxy, NGC 4449, \\citet{2008NGC4449R} found 3-5 Myr old clusters with relatively low extinctions ($ A_V \\leq 1.5$ mag). Although these very early stages are fundamental for our understanding of cluster formation, evolution and possible destruction (90 \\% of clusters are not expected to survive after their first 10 Myr, \\citealp{2003ARA&A..41...57L}), they are still poorly understood.  The cluster formation process is intimately connected both to the galactic environment via the star formation efficiency (SFE), and on long time scales, to the global star formation history (SFH) of the galaxy. Interactions and/or merging between gas-rich galaxies provide physical conditions which favour the formation of star clusters and enhance star formation rates (SFR). Star clusters are easily detected in active star forming regions inside normal spiral galaxies (\\citealp{2002AJ....124.1393L}; \\citealp{2009A&A...503...87C}); in starburst galaxies, like M\\,82 (\\citealp{2009ApJ...701.1015K} and references therein) and NGC\\,3310 (\\citealp{2003MNRAS.342..259D}; \\citealp{2003MNRAS.343.1285D}); interacting \\citep{2005A&A...431..905B} and merging systems (\\citealp{1999AJ....118.1551W}; \\citealp{2008MNRAS.384..886V}; \\citealp{2006ApJ...641..763W}) to cite only a few examples. Interestingly, very bright and massive clusters have also been detected in smaller and less regular systems, like the dwarf post-starburst galaxies NGC 1569 \\citep{2004MNRAS.347...17A} and NGC 1709 (\\citealp{2002AJ....123.1454B}; \\citealp{2009AJ....138..169A}) and in blue compact galaxies (BCGs), which are active dwarf starburst galaxies with low metallicity (\\citealp{2003A&A...408..887O}, \\citealp{adamo09}; Adamo et al.\\, in prep.).  The cluster initial mass function (CIMF) and the cluster luminosity function (CLF) are directly connected to the star formation processes which act from stellar to galactic scales, due to the clustered nature of star formation (\\citealp{2003ARA&A..41...57L}). In general, the CIMF can be approximated by a power law relation of the type $dN(M)/dM=CM^{-\\gamma}$, where $\\gamma \\approx 2$ (\\citealp{1999ApJ...527L..81Z}; \\citealp{2003A&A...397..473B}; \\citealp{2003AJ....126.1836H}, etc.). The CLF, $dN(L) \\propto L^{-\\alpha} dL$, on the other hand, has been shown to have wider range of the index values, $1.8 \\leq \\alpha \\leq 2.7$, and in some cases is better fit with a double power law with a break in the range $-10.4 \\leq M_V \\leq -8$ mag (\\citealp{2002AJ....124.1393L}; \\citealp{1999AJ....118.1551W}; \\citealp{2006A&A...446L...9G}). In a galaxy, the CLF is the synthesis of several populations of SCs assembled at different ages. For this reason, the CLF contains imprints of both the mechanisms behind cluster evolution and the initial environmental conditions where the clusters form. A number of studies have shown that high star formation levels enable a better sampling of the LF up to its brighter end and in turn determine the formation of the brightest cluster in a galaxy (\\citealp{2002AJ....124.1393L}; Gieles et al. 2006a,b; \\citealp{2008MNRAS.390..759B}; \\citealp{2009A&A...494..539L}; \\citealp{2009MNRAS.394.2113G}). On the other hand, physical factors also put constraints on the mass (and consequently the luminosity) of the clusters which are born, such as mass of the parent giant molecular cloud and the pressure and density of the interstellar medium (\\citealp{2003ARA&A..41...57L} ; \\citealp{2002AJ....123.1454B} , \\citealp{2003AJ....126.1836H}). %It is not yet clear how physical %processes and statistics compete in shaping the observed CLF. \\citet{2006A&A...450..129G} noted that if physical conditions in the galactic environment limit the maximum value of the cluster mass, leading to a truncated CIMF, this effect will appear in the LF as a break which corresponds to the maximum cluster mass at the oldest age the galaxy can form.  In this paper we focus our attention on the blue compact galaxy Haro 11 (ESO\\,350-IG038), which has been shown to be extremely efficient in producing SCs. Generally, BCGs have magnitude between M$_B \\approx -12$ and M$_B \\approx -21$ \\citep{2000A&ARv..10....1K}. They are characterized by an ongoing active starburst phase with high star formation rates (SFRs), producing their typical high surface brightness and prominent nebular emission line spectra. Photometric studies of their galactic stellar populations have revealed the presence of an old underlining stellar component (\\citealp{1988A&A...204...10K}; \\citealp{2002A&A...390..891B}; \\citealp{genoveva}). Thanks to the {\\it HST}'s resolution power, it has been possible to resolve the BCG burst regions into many bright star clusters implying that the star formation is remarkably efficient. However, it still not clear yet how systems with masses $\\leq 10^{10} \\msun$ can form such large numbers of star clusters. \\citet{2002AJ....123.1454B}, investigating a sample of 22 nearby ($< 7$ Mpc) dwarf starburst and post-starburst galaxies, noticed that dwarf starburst galaxies have luminosities brighter than $M_V=-16$ and pointed out that despite their sizes, they are able to host very massive clusters. Whether this is a consequence of stochastic sampling effects of the CIMF in combination with environmental conditions which favour the formation of star clusters with $M \\geq 10^5 \\msun$ or different cluster formation physics is not yet established. In this paper, we will call clusters with masses $\\geq 10^5 \\msun$ super star clusters (SSCs).  In Figure~\\ref{h11}, we show a three-colour image of Haro 11. The three main starburst regions, called knots $A$, $B$, and $C$ by \\citet{2003ApJ...597..263K}, are easily traced in the heart-shape. Dust filaments cross knot $B$ and are present in all the burst regions.  Haro 11 is a Ly$\\alpha$ and Lyman continuum emitter \\citep{b1} and can be considered a local analog \\citep{2008ApJ...677...37O} of the high redshift Lyman break galaxies (LBGs). Moreover, Haro 11's IR luminosity of $1.9\\times10^{11}$ L$\\odot$ (section 6.2) locates Haro 11 in the range of the luminous IR galaxies (LIRGs). \\cite{2001A&A...374..800O}, studying the H$\\alpha$ velocity field, found that the mass estimate from rotation curves is lower than the photometric one, probably because the galaxy is not rotationally supported, perhaps because the gas is not in dynamical equilibrium. The total estimated stellar mass is $\\approx 10^{10} \\msun$. Haro 11's perturbed morphology, multiple-component H$\\alpha$ velocity field and high SFR are all signatures of a merger between a low-mass, evolved system and a gas-rich component. Interestingly, \\citet{2001A&A...374..800O} have estimated that the present burst in Haro 11 has been active for $\\sim 35$ Myr, and using the mass of neutral hydrogen, H\\,{\\sc i} estimated by \\citet{2000A&A...359...41B}, a gas consumption time scale (that is, remaining lifetime) of $\\sim$ 5 Myr for the present burst phase. We estimated in section \\ref{cfr_sfr} that the present SFR in Haro 11 is of $22\\pm3 \\msun yr^{-1}$. If we take into account the total mass of gas in H\\,{\\sc i}, H\\,{\\sc ii}, H$_2$ and photodissociation regions, $M_H^{tot} \\sim 2\\times10^9 \\msun$  \\citep{2000A&A...359...41B} we find that the present burst would last at least 100 Myr at its current rate. These time scales show that the gas consumption in the galaxy is currently very rapid but its future development is highly uncertain.  Finally, the similarity between this galaxy and the high-redshift LBGs provides an opportunity to constrain the star-forming mechanisms of the earliest galaxies, when the metallicity was low and the environments were unrelaxed and chaotic.  This work, together with the analysis done by \\citet{2003A&A...408..887O} and two future publications on the BCGs ESO\\,185-IG03 and Mrk\\,930, will be part of a statistical analysis to constrain and define the major properties of the SSC populations in BCGs.  The paper is organized as follows. In section 2 and 3 we describe in detail our data reduction and cluster detection method. Section 4 focuses on the estimation of the clusters' physical properties, in particular ages, masses and extinction. In this section we also assess  the evidence for flux excesses at red wavelengths. Section 5 we examine the possible causes of the red excess. In section 6 we report on the properties of the host galaxy as showed by the cluster analysis. In section 7 we collect our main conclusions.  The luminosity distance and distance modulus of Haro 11 are estimated to be 82.3 Mpc and (m-M)=34.58 mag\\footnote{value taken from NASA/IPAC Extragalactic Database (NED)}, respectively, assuming a cosmology of $H_0 = 73$ Km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M = 0.27$, and $\\Omega_{\\Lambda} = 0.73$ throughout.  \\section[]{Data sample and reductions}  A summary of the full data set used in the present analysis\\footnote{Based on observations made with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555.} is given in Table \\ref{table-obs}.  We have used archival\\footnote{Associated with programs \\# GO 9470 (PI D. Kunth) and \\# GO 10575 (PI G. \\\"Ostlin).} {\\it HST/ACS} imaging of Haro 11, with observations from the FUV to the optical ($\\sim 8000 $\\AA) wavebands. We have supplemented these data with new \\footnote{Associated with programs \\# GO 10902, PI: G. \\\"Ostlin.} {\\it HST/WFPC2} broadband imaging in V (F606W), I (F814W) and H ({\\it NIC 3/F160W}; NICMOS near infrared NIR). To sample further into the infrared we have used the NaCo adaptive optics (AO) imager in the K$_s$ band\\footnote{Based on observations made with ESO Telescopes at the Paranal Observatory under programme IDs 079.B-0585 and 081.B-0234. The observations were carried out the July 24$^{th}$, 2008.}. Unless otherwise specified, we use the Vega magnitude system throughout the paper. The zero points in Vega magnitudes are given in Table \\ref{table-obs} for all the filters.  \\begin{table*} \\caption{Haro 11: {\\it HST} observations. Archival observations are associated with programs \\# GO 9470 (PI D. Kunth) and \\# GO 10575 (PI G. \\\"Ostlin). The new observations were obtained in the program \\# GO 10902 (PI G. \\\"Ostlin). The zero point in Vega magnitude system, and the aperture correction, a$_c$ to the point source photometry are included for each filter. The last two columns show the number of objects detected with a $S/N\\geq5$, and the corresponding magnitude limit in each frame. (a) Shallower image used only to do photometry in the two central cluster regions saturated in the deeper exposure.} \\centering \\begin{tabular}{|c|c|c|c|c|c|c|}%{@{}llll|@{}} \\hline Filter &Instrument &  Exposure time & ZP (mag)&a$_c (mag)$ &N($\\sigma\\leq 0.2$)&  mag limit\\\\ \\hline &Archival obs.&&\\\\ \\hline F140LP (FUV)&ACS/SBC&2700 s&20.92&-0.49$\\pm$0.05&113&25.1\\\\ F220W (NUV)&ACS/HRC&1513 s&21.88&-0.41$\\pm$0.05&80&23.2\\\\ F330W (U)&ACS/HRC&800 s&22.91&-0.35$\\pm$0.05&90&23.7\\\\ F435W (B)&ACS/WFC&680 s&25.79&-0.50$\\pm$0.05&178&27.5\\\\ F550M (V) &ACS/WFC&471 s&24.88&-0.45$\\pm$0.05&144&26.2\\\\ F814W (I)&ACS/HRC&100 s&24.86&-0.56$\\pm$0.07&(a)&(a)\\\\ \\hline &New obs.&&&\\\\ \\hline F606W (R)&WFPC2/PC&4000 s&22.89&-0.61$\\pm$0.09&211&27.5\\\\ F814W (I)&WFPC2/PC&5000 s&21.64&-0.73$\\pm$0.04&211&26.2\\\\ F160W (H)&NIC 3&5000 s&21.88&-2.45$\\pm$0.37&71&25.6\\\\ K$_s$&VLT/NaCo&1222 s&23.32&-2.15$\\pm$0.64&66&20.5\\\\ \\hline \\end{tabular} \\label{table-obs} \\end{table*}   \\subsection{{\\it ACS} data reduction} A full description of the data reduction of the archival data is given in \\cite{2009AJ....138..923O}.  The final ACS images have all been drizzled and aligned using the {\\tt MULTIDRIZZLE} task (\\cite{2002hstc.conf..337K}; \\cite{2002PASP..114..144F}) in {\\tt PyRAF/STSDAS}\\footnote{STSDAS and PyRAF are products of the Space Telescope Science Institute, which is operated by AURA for NASA} and rescaled to the HRC pixel scale, that is $0.025 \"/px$. The final science images are rotated to North-up. %The zero points (ZP) have been %calculated using the available {\\tt STSDAS/SYNPHOT} task. The short {\\it HRC} F814W exposure was only used for estimating the flux of the two brightest sources, in the center of knots $B$ and $C$ (Figure~\\ref{h11}). Both knot centers were saturated in the deep {\\it WFPC2} F606W and F814W exposures.  \\begin{figure*} \\resizebox{0.8\\hsize}{!}{\\rotatebox{0}{\\includegraphics{haro11_px.ps}}} \\caption{The starburst regions in Haro 11. The image is the result of the combination of the {\\it WFPC2} F814W filter in red, and the two {\\it ACS} filters: F435W in green and F220W in blue. The orientation and the $1\"$ length are indicated. The names of the three starburst regions are the same as in \\citet{2003ApJ...597..263K}.} \\label{h11} \\end{figure*} \\subsection{{\\it WFPC2} data reduction} Due to the small size (major axes $\\sim 13.0\"$) of the starburst regions into the galaxy, the target could easily fit into the high-resolution (pixel scale is $0.045\"$) Planetary Camera (PC) which has a FOV of $35\"\\times35\"$. In order to refine the WFPC2 resolution we used sub-pixel dithering and combining patterns. The data reduction were performed with {\\tt MULTIDRIZZLE}. %In this way, setting {\\it %ad-hoc} parameters, we could correct the images for geometric %distortions, bad pixels, and cosmic rays. We took into account sky corrections during the photometric analysis rather than performing sky-subtraction with {\\tt MULTIDRIZZLE}. Before combining the final image, we resized the pixel scale to $0.025\"$, applying a drop size of $0.8$ for both the filters, and rotated the final image so that North was up.  \\subsubsection{The {\\it WFPC2} point spread function}  Because we found no isolated bright objects either in the galaxy or in the PC field, we constructed a filter-dependent point spread function (PSF) using archival data. We retrieved {\\it WFPC2}/F814W images of the NGC-104-F field of the globular cluster 47 Tuc (\\# GO 8160, PI I.\\ King). This field was chosen because it is not crowded and has dithering path and exposure time comparable to our data.  We built the F606W PSF in a similar fashione using data from another uncrowded stellar field (\\# GO 8090, PI S. Casertano). We reduced the images applying the same method, taking into account only the PC field of the mosaic. Isolated sources in the field were selected by eye, and encircled energy distributions for each of them were inspected in order to avoid sources with irregular PSF shapes. %We carried out %photometry on the selected candidates with the {\\tt PyRAF} task {\\tt %PHOT} with an aperture radius of $0.45\"$ and sky annulus at %$0.5-0.6\"$. The tasks {\\tt PSTSELECT} and {\\tt PSF} were used to produce the final PSF. These two constructed PSFs were used to perform a completeness test (see section 3.2).  \\subsection{{\\it NICMOS} data reduction}  The NIC3 camera was chosen because of the size of the field of view, $51.2\" \\times 51.2\"$. To sub-sample and improve the pixel scale resolution (for NIC3, the plate scale is $\\sim 0.2\"/px$) a double NIC-SPIRAL-DITH path with 3 dithered pointings each was applied in order. The single science frames showed significant differences in the sky background levels and presence of bias residuals, which we corrected before running {\\tt MULTIDRIZZLE}. In order to get the best combination of pixel sub-sampling and pixel drop size values we compared the flux growth curves of some bright point-like objects in the final science frame with a reference model PSF produced by {\\tt TinyTim} software\\footnote{http://www.stsci.edu/software/tinytim/}. The resulting best set of input values for the final combined image was found to be {\\tt pixscale} $= 0.067\"/px$ and {\\tt pixfrac} $=0.7$. The pixel scale applied to the final frame is the diffraction limited value for the NIC3 camera. The final science frame was rotated to North-up and calibrated in units of $DN/s$. %To estimate the ZP in Vega %magnitude, we used the system tables published on the {\\it NICMOS} web %page and the relations in Chapter 5 of the {\\it NICMOS Data %Handbook}. We performed a number of tests of the calibration of the final IR frame: these are detailed in Appendix A.  \\subsection{NaCo data reduction}  Ks imaging was carried out at the VLT UT4 telescope using Naos-Conica (NaCo) AO instrument in natural guide star mode. The CONICA imager and spectrograph equipped with the Aladdin3 array and S27 camera (FOV of $28\"\\times28\"$) was chosen. The chopping mode and dithering path were set and a total of 24 target images and 24 sky fields were taken, each with $DIT = 9.4 s$ averaged on 13 $NDITs$. We used {\\tt PyPELINE} \\citep{genoveva} for bad-pixel masking, pair subtraction, flat field correction, source masking, sky background subtraction and alignment. Before the final stacking, the quality of every frame was checked and those with blurred PSFs were excluded, in order to achieve the best possible resolution. Frames with PSF full width at half maximum (FWHM) $> 4 px$($\\sim 0.108\"$) were rejected. The final image was built from the median of the best 10 selected frames (corresponding to a total exposure time of 1222 s). A standard star was observed the same night, with the same instrumental set, and at comparable airmass values that for our target (the difference is $< 0.08$) so that the ZP value is affected by the same atmospheric extinction value as the galaxy. The pixel scale in the final frame was $0.027\"/px$. ",
        "conclusions": "Thanks to the deep {\\it HST} data stretching from UV to IR wavebands, we have carried out a photometric study of the clusters in the starburst regions of Haro 11. We have constrained the physical properties of 185 clusters through the analysis of their integrated SEDs. We summarize the main results here. \\begin{enumerate} \\item More than 60 \\% of the star clusters in Haro 11 are younger than 10 Myr old. The present starburst phase started less than 40 Myr ago with a peak as recently as 3.5 Myr. Only 16 clusters have estimated ages greater than 40 Myr and, of these, half are older than $> 100$ Myr. Haro 11 therefore provides a unique opportunity to investigate the cluster formation and evolution in a starburst system in its earliest phases.  For ages between 1 and 3.5 Myr the extinction range is quite wide, consistent with a complex cluster formation process and the gradual dissolution of the parent clouds. %Moreover, we observed that extinction in general tends to %decrease with cluster age. \\item The clusters are quite massive, with masses between $10^3 - 10^6 \\msun$. There are around 60 estimated SSCs, 7 of them with masses larger than $10^6 \\msun$, and two very massive knots (B and C) with estimated masses $\\geq 10^7 \\msun$ which we could not resolve into smaller clusters. Very young clusters (1-3 Myr) with masses around $10^3-10^4 \\msun$ are absent from our data even though our detection limits should have allowed their detection. These missing clusters could still be embedded in their parental clouds, invisible in optical wavebands, or blended with other clusters in the most crowded regions. \\item We have discovered that a fraction of the clusters exhibits a flux excess compared to plausible SED models at wavelengths $>8000$ \\AA.  We have noted that fitting those cluster SEDs affected by the red excess with normal SSP models leads to overestimation of the clusters ages and masses. %Moreover, attempting fits to the %NIR bands for these clusters restricts models to older %ages for which lower-mass stars dominate the light. In several cases, we have obtained acceptable fits  only at bands blueward of I,  highlighting the potential impact of the phenomenon  on  estimations of the cluster properties. Of the  clusters affected by the red excess, none are older than 40 Myr. \\item We have discussed possible causes for the red excess. If the age of the cluster is the key property, then the red excess in clusters younger than 5 Myr could be due to the simultaneous presence of very young UV-bright massive stars and the dust left from the GMCs. The ERE mechanism could then originate the excess at I waveband while hot dust emission produces a \"rise\" at IR wavelengths, depending on the size and temperature of the dust grains. Another possible mechanism at such young stages could be the presence in the cluster of two stellar components, one of which has shed its dust envelope, contributing mainly in the UV-optical wavelengths, while the other component is still deeply embedded in its star-forming cloud and highly extinguished. Pre-main sequence objects and young stellar objects (YSOs) still surrounded by circumstellar disks can produce this IR-bright component. Further modelling will be required in order to estimate how massive the second stellar component needs to be in order to produce such large displacements in the H and K$_s$ bands.  The presence of such embedded stellar populations emitting only in the IR would also imply that the mass of these clusters will be underestimated. At ages $> 5 Myr$ we expect that the embedded phase is over; we have suggested two further mechanisms to explain the NIR excess for ages older than this. First, evolved clusters may be observed in the line of sight towards younger, embedded ones. In such cases, the cluster behind would only be visible at IR-bands, while the age and mass estimated form UV and optical bands correspond to the cluster in front. Second, red supergiants, which dominate the light at red and near-IR wavelengths, could \"redden\" the observed colours of the clusters at ages as early as 6 Myr. %Finally, if the age of the cluster is not after all the key %issue, a bottom-heavy IMF could %in principle explain the IR excess. If this is the case, it is  difficult to explain why only a fraction of the clusters are producing stars with a %different IMF, and the the total stellar mass necessary to produce %such an excess is uncomfortably high. \\item For the first time, after an early attempt by \\cite{2003A&A...408..887O}, we have analysed the cluster luminosity function in a BCG. Haro 11's CLF is shallower than observed in other galaxies. Contamination by blending could have this effect on the CLF, but we didn't find any evidence of overestimation of the CFR, using the more massive SSCs. Obviously, at the distance of the galaxy and with the present available data, we cannot avoid crowding and projection contaminations. \\item More than the 30\\% of the star formation in Haro 11 has happened in bound clusters which can survive and evolve along with the galaxy. This fraction is higher than observed in the Milky Way and other merging galaxies like NGC 3256, leading to the conclusion that the environmental conditions in this BCG have favoured the formation of massive super star clusters. Due to the similarities between Haro 11 and the high-redshift LBGs, we suggest that cluster formation rates at higher redshift may also be enhanced. \\end{enumerate}"
    },
    "1005/1005.2234_arXiv.txt": {
        "abstract": "We examine the relation between black hole accretion and bulge star formation as a function of look-back time ($\\tau$) in 20,541 obscured AGNs (with redshifts $\\bar{z}\\sim$0.1 and bolometric luminosities $L_{{\\rm Bol}}\\sim10^{43}$--$10^{45}$ erg s$^{-1}$) optically selected from the Sloan Digital Sky Survey (SDSS). To quantify the most recently formed stars with ages less than typical AGN lifetimes, we estimate the differentiated specific star formation rate (SSFR$_{\\tau}$) based on population synthesis analysis.  Eddington ratio ($\\lambda$) is inferred using \\OIIIb\\ luminosity and stellar velocity dispersion as proxies for $L_{{\\rm Bol}}$ and black hole mass respectively. We find that when $\\tau < \\tau_{0}$, SDSS AGNs follow a power law $\\lambda \\propto$ SSFR$_{\\tau}^{1.0-1.1}$; the relation flattens out when $\\tau > \\tau_{0}$. The threshold timescale $\\tau_{0}$ is $\\sim 0.1$ ($\\sim 1$) Gyr in young (old) bulges. The scatter in the power laws is dominated by observational uncertainties. These results may provide useful constraints on models explaining the correlations between AGN activity and bulge star formation. ",
        "introduction": "\\label{sec:intro}  The growth of supermassive black holes (SMBHs) could be linked with star formation (SF) in their host-galaxy bulges through multiple processes \\citep[see e.g.,][for a review]{heckman08}. Stars form under the cosmological infall of gas \\citep{binney77,rees77}, some of which could be accreted by central SMBHs \\citep[e.g.,][]{burkert01}.  Galaxy mergers or interactions may trigger starburst and quasar activity \\citep[e.g.,][]{sanders88,moore96}, which could output energy and momentum regulating further growth \\citep[e.g.,][]{silk98,fabian99,springel05}.  AGN-driven outflows could boost starbursts \\citep[e.g.,][]{begelman89,silk09}. After being active, evolved stars may fuel AGNs via mass losses \\citep[e.g.,][]{norman88,ciotti97}.  These processes may have different significance in regulating SMBH growth and bulge SF in various AGN populations at different cosmic epochs.   While post-starburst populations have been extensively observed in quasars \\citep[e.g.,][]{boroson82,canalizo00,jahnke07,liu09} and in higher-luminosity AGNs \\citep[e.g.,][]{kauffmann03,heckman04,cid04,netzer09}, the physical connection is still in debate.  Two useful parameters are specific star formation rate (${\\rm SSFR} \\equiv {\\rm SFR}/M_{\\ast}$) and Eddington ratio ($\\lambda \\equiv L_{{\\rm Bol}}/L_{{\\rm Edd}}$).  By studying the $\\lambda$ distribution in AGNs as a function of mean stellar age, \\citet{kauffmann09} suggested that local AGNs occupy two distinct growth regimes. These authors used the 4000-{\\rm \\AA} break \\citep[D$_n(4000)$;][]{bruzual83,balogh99} to characterize mean stellar age.  Arising from a series of metallic lines, D$_n(4000)$ is most sensitive to 1--2 Gyr old stars \\citep[e.g.,][]{bc03}, whereas luminous AGN phases appear to be brief \\citep[$<$0.1 Gyr; e.g.,][]{yu02}.  To get higher temporal resolution and to quantify the most recently formed stars with ages less than typical AGN lifetimes, we examine SSFR averaged in the past $\\tau$ years (SSFR$_{\\tau}$) based on population synthesis analysis \\citep[e.g.,][]{bica88,cid05,asari07}.  Under the multiple-starburst approximation, the differentiated SF histories in a statistical sample of AGNs offered by the Sloan Digital Sky Survey \\citep[SDSS;][]{york00} enable us to investigate the correlation between AGN and recent bulge SF activity as a function of look-back time.  This approach offers insights otherwise unattainable with single and older age indicators.  We describe the AGN sample and the assumptions in estimating $\\lambda$ and SSFR$_{\\tau}$ along with uncertainties and selection biases in \\S \\ref{sec:sample}. \\S \\ref{sec:corr} focuses on the correlations between $\\lambda$ and SSFR$_{\\tau}$ as a function of $\\tau$.  We discuss implications and caveats in \\S \\ref{sec:discuss}.  A cosmology with $\\Omega_m = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $h = 0.7$ is assumed throughout. ",
        "conclusions": "\\label{sec:discuss}  Our results imply similar relations between AGN activity and recent SF in old and young bulges, despite different Eddington ratios observed.  The current Eddington ratio correlates with the most recent bulge SF: $\\lambda \\propto$ \\ssfr\\ estimated in a certain look-back time interval $\\tau$; it depends little on SF happened before some threshold time $\\tau_{0}$.  The main difference between young and old bulges is on $\\tau_{0}$: young bulges have $\\sim 0.1$ Gyr whereas old bulges have $\\sim 1$ Gyr.  Recently \\citet{kauffmann09} have studied the $\\lambda$ distribution as a function of D$_n$(4000). These authors find that P($\\lambda$) is log normal in young bulges and exponential in old bulges, and suggest two growth regimes.  While this could still be the case, \\citet{hopkins09} suggest that different $\\lambda$ distributions may also arise from universal AGN light curves with different quench times: AGNs in old bulges have apparent exponential P($\\lambda$) because accretion and star formation were quenched in a series of earlier epochs \\citep[see also e.g.,][]{shen09}.  Our results point to the latter explanation.  Namely, the physical processes that link AGN activity and bulge SF seem to work universally in young and old bulges, albeit with different quenching/threshold times.  Using population synthesis analysis, we have found that when grouped in terms of similar SF activity with ages less than typical AGN lifetimes, BHs in old bulges also have log normal P($\\lambda$).  We have shown that for BHs in young bulges, $\\lambda$ does correlate with the most recent SF (Figs. \\ref{fig:ssfr}a,b).  The reason why \\citet{kauffmann09} found little correlation is most likely that the temporal resolution of D$_n$(4000) is not enough to disintegrate the most recent SF in young bulges. The reciprocal SSFR$_{\\tau}^{-1}$, which is the mass-weighted stellar age averaged among the stellar populations formed over $\\tau$, can be estimated for different timescales $\\tau$, whereas D$_n$(4000) is most sensitive to $\\tau = $ 1--2 Gyr.  If we assume the observed sample consists of AGNs at different evolutionary stages, the $\\lambda$-SSFR results (Table \\ref{table:fit}) imply light curves $\\lambda \\propto t^{-(1.0-1.1)}$, which is noticeably less steep than predictions $\\lambda \\propto t^{-(1.5-2.0)}$ from models assuming extreme gas expel by AGNs \\citep[e.g.,][]{hopkins09}. This milder decay could suggest more moderate self-regulated AGN and SF activity, although the observational uncertainties in our results are significant.   Another implication is that AGN environmental studies should draw control samples matched not only in the mean stellar age (as usually implemented using D$_n$(4000)), but also in the most recent bulge SF properties.  The effect would be particularly important for AGNs in young bulges, as we have shown that their SF-AGN-correlation timescale is usually much shorter ($\\sim$ 0.1 Gyr) than that probed by the overall mean stellar age.   Our results are limited to optical AGNs with $\\bar{z}\\sim$0.1 and $L_{{\\rm Bol}}\\sim10^{43}$--$10^{45}$ erg s$^{-1}$, and may not apply to other AGN populations.  In principle the approach could be extended to lower-luminosity AGNs (LINERs) and to higher-luminosity AGNs and quasars, although selection incompleteness and biases may be more severe. Multi-wavelength SFR indicators that probe different SF timescales or phases may also help elucidate the AGN-SF relation as a function of look-back time.  The approach demonstrated based on population synthesis analysis relies on statistical AGN samples with high-S/N spectra. We look forward to the next generation of ground-based near-IR multi-object spectrographs \\citep[e.g.,][]{gunn09} that will make an equivalent study possible at high redshifts to better understand the coupled growths of stellar bulges and SMBHs in the early universe."
    },
    "1005/1005.5552_arXiv.txt": {
        "abstract": "We present here the final results of the first spectropolarimetric survey of a small sample of active M dwarfs, aimed at providing observational constraints on dynamo action on both sides of the full-convection threshold (spectral type M4). Our two previous studies \\cite[][]{Donati08b, Morin08b} were focused on early and mid M dwarfs. The present paper examines 11 fully convective late M dwarfs (spectral types M5--M8). Tomographic imaging techniques were applied to time-series of circularly polarised profiles of 6 stars, in order to infer their large-scale magnetic topologies. For 3 other stars we could not produce such magnetic maps, because of low variability of the Stokes~V signatures, but were able to derive some properties of the magnetic fields.  We find 2 distinct categories of magnetic topologies: on the one hand strong axisymmetric dipolar fields (similar to mid M dwarfs), and on the other hand weak fields generally featuring a significant non-axisymmetric component, and sometimes a significant toroidal one. Comparison with unsigned magnetic fluxes demonstrates that the second category of magnetic fields shows less organization (less energy in the large scales), similarly to partly convective early M dwarfs. Stars in both categories have similar stellar parameters, our data do not evidence a separation between these 2 categories in the mass-rotation plane.  We also report marginal detection of a large-scale magnetic field on the M8 star VB~10 featuring a significant toroidal axisymmetric component, whereas no field is detectable on VB~8 (M7). ",
        "introduction": "\\label{sec:intro} Magnetic field is a key parameter in theories of stellar formation and evolution. In cool stars, it powers several activity phenomena, observed on a wide range of wavelengths and timescales, which provide a rough proxy of the averaged field strength. The well-established rotation-activity relation \\cite[e.g., ][]{Noyes84} supports the idea of a dynamo-generated magnetic field in these stars.  \\cite{Larmor19} first proposed that the solar magnetic field could be induced by plasma motions, already pointing out the importance of shear ($\\Omega$ effect) to generate strong toroidal fields and thus explain Hale's polarity law of sunspots. \\cite{Parker55} completed the basic picture of the solar dynamo by introducing the $\\alpha$ effect (convection made cyclonic by the Coriolis force) to address the issues of Cowling's anti-dynamo theorem \\citep{Cowling34} and the regeneration of a poloidal field component from a toroidal one. The $\\alpha\\Omega$ dynamo has thereafter been thoroughly debated and improved \\citep[e.g. ][]{Babcock61, Leighton69}. More recently, helioseismology has been able to probe the solar interior and revealed the existence of a thin layer of strong shear located at the base of the convection zone: the tachocline \\cite[e.g.,][]{Spiegel92}. Although many aspects of the solar magnetism are still not thoroughly understood, recent theoretical and numerical studies have pointed out the crucial role of the tachocline as the place of storage and amplification of strong toroidal fields \\cite[e.g., ][]{Ossendrijver03, Charbonneau05}.  New insight on dynamo processes can be gained from the exploration of magnetic fields of cool stars, probing the dynamo response in very non-solar regimes of parameters (e.g., fast rotation, deeper or shallower convection zone). M dwarfs are particularly interesting since those below $\\sim 0.35~\\msun$ \\cite[e.g.,][]{Chabrier97} are fully convective and therefore do not possess a tachocline and presumably cannot host a solar-type dynamo. Yet, many M dwarfs are known to be active, and these stars follow the usual rotation-activity relation \\cite[e.g., ][]{Delfosse98, Reiners07, Kiraga07, West04}. However the strong correlation between X-ray and radio luminosities established by \\cite{Guedel93} for stars of spectral types ranging form F to mid M, is no longer valid for very low mass dwarfs which exhibit very strong radio emission whereas X-ray emissions dramatically drop \\cite[][]{Berger06}. Magnetic fields were also directly detected at photospheric level through Zeeman effect both in unpolarised \\cite[e.g., ][]{Saar85, Johns96, Reiners06} and circularly polarised \\citep{Donati06} line profiles.  Spectropolarimetry combined with tomographic imaging techniques is the optimal technique to investigate the magnetic topologies of M dwarfs (see Sec.~\\ref{sec:model} for more details). By recovering the large-scale component of stellar magnetic fields, we can provide dynamo theorists with observables directly comparable with their modeling (axisymmetry, relative importance of the poloidal and toroidal components, characteristic scales...). Previous results have already produced strong constraints: \\cite{Donati06} and \\cite{Morin08a} (hereafter M08a) demonstrated that the fully convective fast rotator V374~Peg is able to trigger a strong large-scale axisymmetric poloidal field steady on a timescale of 1~year; and exhibits a very low level of differential rotation ($\\sim \\frac{\\dOmsun}{10}$).  From the analysis of a sample of early and mid M dwarfs  \\cite{Donati08b} and \\cite{Morin08b} (hereafter D08 and M08b respectively) observed a strong change near the theoretical full-convection threshold: while partly convective stars possess a weak non-axisymmetric field with a significant toroidal component, fully convective ones exhibit strong poloidal axisymmetric dipole-like topologies. Differential rotation also drops by an order of magnitude across the boundary, the observed fully convective stars exhibit nearly solid body rotation. D08 and M08b also report a sharp transition in the rotation-large-scale magnetic field relation close to the full-convection boundary, whereas no such gap is visible in the rotation-X-ray relation. Considering that X-ray emission is a good proxy for the total magnetic energy, it suggests that a sharp transition in the characteristic scales of the magnetic field occurs near the full-convection limit. This point was further confirmed for a few stars by \\cite{Reiners09b} who report that the ratio of the magnetic fluxes measured from circularly-polarised and unpolarised lines dramatically changes across the fully convective limit. Several theoretical studies have addressed the challenging issue of dynamo action in fully convective stellar interiors. \\cite{Durney93} first proposed that without a tachocline of shear, convection and turbulence should play the main role at the expanse of differential rotation, generating a small-scale field. \\cite{Kuker99} and \\cite{Chabrier06} performed mean-field modeling of dynamo action in fully convective stars and found purely non-axisymmetric $\\alpha^2$ solutions, indicating that these objects can sustain large-scale magnetic fields. Subsequent direct numerical simulations by \\cite{Dobler06} and \\cite{Browning08} both realized large-scale dynamo action with a significant axisymmetric component of the resulting magnetic field. The latter also achieved magnetic energy in equipartition with kinetic energy and therefore Maxwell stresses strong enough to quench differential rotation, resulting in nearly solid-body rotation. Despite these recent advances, the precise causes of differences between dynamo in fully and partly convective stars are not completely understood, and theoretical studies need observational guidance.  In the present paper, we extend our spectropolarimetric study to 11 late M dwarfs (spectral types ranging from M5 to M8). After a brief presentation of the stellar sample and of spectropolarimetric observations, we describe the main principles of the tomographic imaging process. The Zeeman-Doppler Imaging (ZDI hereafter) analysis is then detailed for 6 stars. For 3 other late M dwarfs, it is not possible to derive a definitive magnetic map because of the very low level of variability in the Stokes~$V$ signatures, but we can still infer some information about their magnetic topologies. For the very faint stars  VB~8 and VB~10, the noise level is too high to allow definite detection of the circularly polarised signatures in individual LSD spectra. By averaging the spectra of each data set, we marginally detect a large-scale magnetic field on VB~10, and show a tentative ZDI reconstruction. We finally discuss these results and conclude on the implications of our study for the understanding of dynamo processes in fully convective stars. ",
        "conclusions": "\\label{sec:disc} \\begin{table*} \\caption[]{Magnetic quantities derived from our study. For each star, different observation epochs are presented separately. In columns 2--4 we report quantities from Table~\\ref{tab:sample}, respectively the stellar mass, the rotation period, and the effective Rossby number. Columns 5, 6 and 7 mention the Stokes $V$ filling factor, the reconstructed magnetic energy and the average magnetic flux. Columns 8--11 list the percentage of reconstructed magnetic energy respectively lying in poloidal, dipole (poloidal and $\\ell=1$), quadrupole (poloidal and $\\ell=2$), and octupole (poloidal and $\\ell=3$) modes. In column 12, we mention the percentage of magnetic energy reconstructed in axisymmetric modes (defined as $m < \\ell/2$) and the percentage of poloidal energy in axisymmetric modes. See section~\\ref{sec:techniques-uncert} for a discussion on the robustness of magnetic map reconstruction and the uncertainties associated with the derived quantities. } \\begin{tabular}{ccccccccccccc} \\hline Name & Mass & \\Prot & $Ro$ & $f_V$ & $\\avg{B^2}$ & $\\avg{B}$ & pol. & dipole & quad. & oct. & axisymm. \\\\ & (\\msun) & (d) & ($10^{-2}$) & & ($\\rm10^5\\,G^2$) & (kG) & (\\%) & (\\%) & (\\%) & (\\%) & (\\%) \\\\ \\hline GJ~51$^1$ (06) & 0.21 & 1.02 & 1.2 & 0.12 & 38.6 & 1.61 & 99 & 96 & 0 & 2 & 91/91\\\\ \\phantom{GJ~51} (07) &--&--&--& 0.12 & 31.3 & 1.58 & 99 & 92 & 0 & 6 & 77/77\\\\ \\phantom{GJ~51} (08) &--&--&--& 0.12 & 32.6 & 1.65 & 97 & 92 & 1 & 3 & 89/89\\\\ GJ~1156$^2$ (07) & 0.14 & 0.49 & 0.5 & 1.0 & 0.06 & 0.05 & 88 & 30 & 26 & 19 & 6/3\\\\ \\phantom{GJ~1156} (08) &--&--&--& 1.0 & 0.19 & 0.11 & 83 & 41 & 28 & 11 & 20/12\\\\ \\phantom{GJ~1156} (09) &--&--&--& 1.0 & 0.13 & 0.10 & 94 & 54 & 24 & 10 & 2/1\\\\ GJ~1245~B (06) & 0.12 & 0.71 & 0.7 & 0.06 & 0.44 & 0.17 & 80 & 45 & 14 & 13 & 15/9\\\\ \\phantom{GJ~1245~B} (07) &--&--&--& 0.10 & 0.49& 0.18 & 84 & 46 & 27 & 7 & 52/53\\\\ \\phantom{GJ~1245~B} (08) &--&--&--& 0.10 & 0.06 & 0.06 & 85 & 33 & 25 & 19 & 20/18\\\\ WX~UMa (06) & 0.10 & 0.78 & 0.8 & 0.12 & 16.08 & 0.89 & 97 & 66 & 21 & 6 & 92/92\\\\ \\phantom{WX~UMa} (07) &--&--&--& 0.12 & 24.42 & 0.94 & 97 & 71 & 13 & 3 & 92/94\\\\ \\phantom{WX~UMa} (08) &--&--&--& 0.12 & 23.53 & 1.03 & 97 & 69 & 19 & 6 & 83/85\\\\ \\phantom{WX~UMa} (09) &--&--&--& 0.12 & 37.54 & 1.06 & 96 & 89 & 2 & 2 & 95/96\\\\ DX~Cnc (07) & 0.10 & 0.46 & 0.5 & 0.20 & 0.17 & 0.11 & 93 & 69 & 11 & 9 & 77/77\\\\ \\phantom{DX~Cnc} (08) &--&--&--& 0.20 & 0.09 & 0.08 & 73 & 31 & 25 & 10 & 49/34\\\\ \\phantom{DX~Cnc} (09) &--&--&--& 0.20 & 0.09 & 0.08 & 62 & 42 & 11 & 4 & 70/61\\\\ GJ~3622 (08) & 0.09 & 1.5 & 1.5 & 1.0 & 0.04 & 0.05 & 96 & 90 & 7 & -- & 73/72\\\\ \\phantom{GJ~3622} (09) &--&--&--& 1.0 & 0.05 & 0.06 & 93 & 84 & 9 & -- & 80/78\\\\ \\hline \\label{tab:syn} \\end{tabular} \\begin{flushleft} $^1$ For GJ~51, the imaging process is weakly constrained by our data sets due to poor phase coverage (see Sec.~\\ref{sec:gj51}).\\\\ $^2$ For GJ~1156 the alternative rotation period 0.33~d cannot be definitely excluded (see Sec.~\\ref{sec:gj1156}). In this case, the reconstructed topologies remain   similar for the 3 epochs, and our conclusions are not affected. \\end{flushleft} \\end{table*}  We present the final part of our exploratory spectropolarimetric survey of M dwarfs, following D08 and M08b (respectively concentrating on mid and early M dwarfs) we focus here on the low mass end of our sample. For 6 stars, it is possible to apply ZDI techniques to the time-series of circularly polarised spectra and thus to infer the large-scale component of their magnetic topologies. The properties of the reconstructed topologies of these stars are presented in Table~\\ref{tab:syn}. For the remaining 5 objects, the data sets do not permit such a study, it is however possible to retrieve some constraints about their magnetic properties.  Two stars of the subsample (namely GJ~51 and WX~UMa) exhibit large-scale magnetic fields very similar to those observed by M08b on mid M dwarfs, i.e. very strong, axisymmetric poloidal and nearly dipolar fields, with very little temporal variations. For three stars for which we cannot perform a definitive ZDI reconstruction (GJ~1154~A, GJ~1224 and CN~Leo) our data strongly suggest similar topologies. From the observations of WX~UMa, we demonstrate that the timescale of temporal evolution in the magnetic topologies of these stars can be larger than 3~years, whereas previous observations by M08a and M08b were based on observations spanning only 1~year.  The other stars for which we can reconstruct the large-scale magnetic topologies, are clearly different. They are weaker than those of the first category, and generally feature a significant non-axisymmetric component, plus a significant toroidal component (although the field is always predominantly poloidal). Temporal variability is also noticeable, in particular for GJ~1245~B our data indicate unambiguously that the magnetic field strength has dramatically decreased between our 2007 and 2008 observations. Our conclusions are robust to uncertainties on stellar parameters (\\Prot, $i$ and $\\vsini$). When varying these parameters over the width of their respective error bars (see section~\\ref{sec:techniques-uncert}), the main properties of the reconstructed magnetic topologies remain unchanged.  \\begin{figure*} \\centering \\includegraphics[height=1.0\\textwidth, angle=270]{fig_arxiv/plotMP.ps} \\caption{Properties of the magnetic topologies of our sample of M dwarfs as a function of rotation period and stellar mass. Larger symbols indicate larger magnetic fields while symbol shapes depict the different degrees of axisymmetry of the reconstructed magnetic field (from decagons for purely axisymmetric fields to sharp stars for purely non axisymmetric fields). Colours illustrate the field configuration (dark blue for purely toroidal fields, dark red for purely poloidal fields and intermediate colours for intermediate configurations). Solid lines represent contours of constant Rossby number $Ro=0.1$ and $0.01$ respectively corresponding approximately to the saturation and super-saturation thresholds \\citep[e.g.,][]{Pizzolato03}. The theoretical full-convection limit ($\\mstar \\simeq0.35\\msun$, \\citealt{Chabrier97}) is plotted as a horizontal dashed line, and the approximate limits of the three stellar groups  discussed in the text are represented as horizontal solid lines. Stars with  $\\mstar > 0.45~\\msun$ are from D08, whereas those with $0.25<\\mstar<0.45~\\msun$ are from M08b. For GJ~1245~B symbols corresponding to 2007 and 2008 data sets are superimposed in order to emphasize the variability of this object. % Uncertainties associated with the plotted magnetic quantities are discussed in section~\\ref{sec:techniques-uncert}.} \\label{fig:plotMP} \\end{figure*}  These results are presented in a more visual way in Fig.~\\ref{fig:plotMP}. Previous studies by D08 and M08b have revealed strong evidence that a clear transition occurs at approximately 0.5~\\msun\\, \\ie more or less coincident with the transition to a fully convective internal structure. The situation here is different, we find stars with similar stellar parameters that exhibit radically different magnetic topologies. On Fig.~\\ref{fig:plotMP}, WX~UMa is the only star below 0.2~\\msun\\ to host a mid-M-dwarf-like field. Whereas DX~Cnc and GJ~1245~B are very close to it in the mass-rotation plane they feature fields with very different properties. This observation may be explained in several ways. For instance, another parameter than mass and rotation period, such as stellar age, may play a role. In our sample we indeed notice that most stars below 0.15~\\msun\\ that exhibit a weak complex field belong to a young kinematic population according to \\citealt{Delfosse98} (GJ~1156, DX~Cnc and GJ~3622), whereas those hosting a strong dipolar field (WX~UMa, GJ~1224 and CN~Leo) belong to older kinematic populations (old disk and old/young disk). This hypothesis requires further investigation. One could also imagine, for instance, that the magnetic fields of very low mass stars may switch between two different states over time. This hypothesis is supported by the fact that for one of the stars in the weak and complex field regime (GJ~1245~B) we observe a dramatic variation of the magnetic flux on a timescale of one year which may indicate that the magnetic field of these objects may go through chaotic variations and eventually switch between the two categories of field actually observed. However no such switch has been observed in our sample. We observe 5 objects in the strong dipole field category (GJ~51, GJ~1154~A, GJ~1224, CN~Leo and WX~UMa) and 6 in the weak field category (GJ~1156, GJ~1245~B, DX~Cnc, GJ~3622, VB~8 and VB~10), indicating that stars would spend as much time in both states. No star is observed in an intermediate state, suggesting that a putative transition would be fast. This hypothesis may be investigated through the analysis of long-term radio monitoring, since radio emission would be presumably strongly impacted by such a dramatic change of the stellar magnetic field. Recent observations of ultracool dwarfs reveal a long-term variability of the activity indices \\cite[\\eg][]{Antonova07, Berger10} that may support this view.  \\begin{figure*} \\centering \\includegraphics[angle=270, width=0.50\\textwidth]{fig_arxiv/Ro_Bf1b.ps}% \\includegraphics[angle=270, width=0.50\\textwidth]{fig_arxiv/Ro_Bf2b.ps}% \\caption[]{Magnetic flux as a function of Rossby number. On the left panel magnetic fluxes as measured from Stokes~$V$ spectra and ZDI by M08b and D08 are mentioned as blue hexagons ($\\mstar>0.4~\\msun$) and green squares ($0.2<\\mstar<0.4~\\msun$). Results from this paper ($\\mstar\\leq0.2~\\msun$) are shown as red circles. On the right panel the ratio between the magnetic fluxes as recovered from Stokes~$V$ and Stokes~$I$ measurements (whenever available) are shown. Stokes~$I$ magnetic fluxes are taken from \\cite{Johns00}, \\cite{Reiners07}, \\cite{Reiners09}, and \\cite{Reiners09b} (see Tab.~\\ref{tab:sample}). Measurements of different epochs (whenever available) for Stokes~$V$ are shown connected by a solid line. On each plot the $Ro=10^{-1}$, corresponding to the saturation level, is depicted as a vertical solid black line. Triangles denote upper limits. In the left panel, horizontal lines show the magnetic fluxes corresponding to saturation for stars with $\\mstar>0.4~\\msun$ (blue), and $0.2<\\mstar<0.4~\\msun$ (green). In the right panel they show the mean fraction of magnetic flux detected in Stokes~$V$ spectra.} \\label{fig:RoBsq} \\end{figure*}  On the left panel of Fig.~\\ref{fig:RoBsq}, we plot the reconstructed magnetic flux as a function of the Rossby number. As mentioned in D08, stars with $\\mstar > 0.4~\\msun$ follow the expected rotation--magnetic field connection, with saturation for $Ro \\lesssim 0.2$. Whereas for $\\mstar > 0.4~\\msun$, we only observe objects in the saturated regime, with a significantly stronger saturation magnetic flux. All the stars of the late M subsample have Rossby numbers below $2\\times10^{-2}$ and are thus expected to be in the saturated regime. As in Fig.~\\ref{fig:plotMP} the stars studied here can be divided into two distinct categories. The first one is composed of stars hosting a very strong magnetic field that lie in the saturated part of the rotation-activity relation, similar to mid M dwarfs. We notice that significantly higher magnetic fluxes are reconstructed for WX~UMa and GJ~51 than for mid M dwarfs studied by M08b, although all these stars seem to lie in the saturated dynamo regime. For the second category of late M dwarfs (having a weak complex field), we recover less magnetic flux (even less than for some early M dwarfs studied by D08), in spite of their very low Rossby number. Super-saturation is unlikely to be the explanation since stars of both categories have similar Rossby numbers. As mentioned for Fig.~\\ref{fig:plotMP}, we observe no object in an intermediate state. The aforementioned variability of GJ~1245~B is also prominent in this plot.  On the right panel of Fig.~\\ref{fig:RoBsq}, magnetic flux inferred from Stokes~$V$ measurements are compared to those derived from Stokes~$I$. Again, stars with masses below and above 0.4~\\msun\\ studied by D08 and M08b clearly form two separate groups. In partly-convective stars only a few percents of the magnetic flux measured in $I$ is detectable in $V$ similarly to what is observed for the Sun, while this ratio is close to 15\\% in fully convective ones. This indicates a higher degree of organization of the field in fully convective mid M dwarfs, with more magnetic flux in the spherical harmonics of lowest degree. The ratios of about 30\\% plotted for WX~UMa are upper limits (since the flux based on Stokes~$I$ is a lower limit), it is thus not clear if this star differs from mid M dwarfs in this respect. This high value however indicates a high degree of organization. For DX~Cnc, GJ~1245~B and GJ~1156, the ratio of Stokes~$V$ and $I$ fluxes is closer to the early M dwarfs value. Therefore the magnetic fields of the weak field category of late M dwarfs share similar properties with that of partly convective M dwarfs.  These results on the magnetic topologies of M dwarfs also suggest that dynamo processes in low-mass main sequence stars and pre-main sequence stars may be similar. Indeed the first spectropolarimetric results on young stars show that the fully convective T Tauri star BP~Tau (0.7~\\msun) exhibit a strong large-scale magnetic field \\cite[][]{Donati08a}, whereas the more massive partly convective star V2129~Oph (1.4~\\msun) possesses a weaker and more complex field \\cite[][]{Donati07}. This is reminiscent of the transition we observe at 0.5~\\msun\\ among main sequence M dwarfs. Recent observations, on the fully convective V2247~Oph (0.35~\\msun) reveal a still weaker and more complex magnetic field \\cite[][]{Donati10} which may correspond to the weak field late M dwarf we present here.  We do not detect Stokes~$V$ signatures in individual spectra of the two latest stars of our sample VB~8 and VB~10. However the averaged LSD profile of VB~10 suggests the presence of a toroidal axisymmetric field component on this object. Further observations may confirm this first spectropolarimetric detection on an ultra-cool dwarf.  The first spectropolarimetric survey of M dwarfs has already provided dynamo theorists with strong constraints on the evolution of surface magnetic fields of M dwarfs across the fully convective divide (D08 and M08b). The results presented in this paper on the magnetic topologies of late M dwarfs reveal a new unexpected behaviour below 0.2~\\msun. We interpret the fact that objects with similar stellar parameters host radically different magnetic topologies as a possible evidence for a switch between two dynamo states (either cyclic or chaotic). Finally our observations suggest the presence of a large-scale magnetic field on the M8 dwarf VB~10, featuring a significant toroidal axisymmetric component, whereas the detection of the magnetic field of VB~8 (M7) is not possible from our spectropolarimetric data."
    },
    "1005/1005.5287_arXiv.txt": {
        "abstract": "We perform an extensive analysis of the \\civ\\,$\\lambda1549$ line in three large spectroscopic surveys of quasars. Differing approaches for fitting the \\civ\\ line can be found in the literature, and we compare the most common methods to highlight the relative systematics associated with each. We choose the line fitting procedure that results in a symmetric profile for the \\civ\\ line and gives accurate fits to local emission features around the line, and use this approach to measure the width of the \\civ\\ line in spectra from the SDSS, 2QZ and 2SLAQ surveys.  The results are compared with a previous study of the \\mgii\\,$\\lambda2799$ line in the same sample. We find the \\civ\\ line tends to be broader than the \\mgii\\ line in spectra that have both lines, and the average ratio between the lines is consistent with a simplistic model for a photoionised, virialised and stratified broad-line region. There exists a statistically significant correlation between the widths of the \\civ\\ and \\mgii\\ lines. However, the correlation is weak, and the scatter around a best fit is only marginally less than the full dynamic range of line widths.  Motivated by previous work on the \\mgii\\ line, we examine the dispersion in the distribution of \\civ\\ line widths. We find that the dispersion in \\civ\\ line widths is essentially independent of both redshift and luminosity. This result is in stark contrast to the \\mgii\\ line, which shows a strong luminosity dependence. Furthermore we demonstrate that the low level of dispersion in \\civ\\ line width ($\\sim0.08$\\,dex) is inconsistent with a pure-disk model for the emitting region and use our data to constrain simple models for the broad-line region.  Finally we consider our results in terms of their implications for the the virial technique for estimating black hole masses. The inconsistency between \\mgii\\ and \\civ\\ line widths in single spectra, combined with the differing behaviour of the \\mgii\\ and \\civ\\ line width distributions as a whole, indicates that there must be an inconsistency between \\mgii\\ and \\civ\\ virial mass estimators. Furthermore, the level of intrinsic dispersion in \\mgii\\ and \\civ\\ line widths contributes less dynamic range to virial mass estimates than the error associated with the estimates. The indication is that the line width term in these UV virial mass estimators may be essentially irrelevant with respect to the typical uncertainty on a mass estimate. ",
        "introduction": "This paper describes an extensive analysis of the \\civ\\,$\\lambda1549$ line width distribution in three large spectroscopic samples of quasi-stellar objects (QSOs). \\civ\\ has the highest ionisation potential of any of the strong broad emission lines in QSO spectra and is the best probe of the high-ionisation inner regions of the broad-line region (BLR). Furthermore, like \\hb\\ and \\mgii\\,$\\lambda2799$, the \\civ\\ line is commonly used to calculate super-massive black hole (SMBH) masses for QSOs. The line is observed in optical spectra with redshifts between $\\sim$1.5 and 5 (corresponding to a look-back time of 9 to 12\\,Gyr), and is the only line used to calculate SMBH masses in the highest redshift objects.  The redshift range between $\\sim$1.5 and 5 is of particular importance to quasar astrophysics since it spans the so called `quasar epoch' at $z\\sim2$ to $3$. Before $z\\sim3$ the space density of QSOs has been observed to increase with time (e.g. \\citealt{osm82,fan01,rich06b}), but since $z\\sim2$ it has fallen (e.g. \\citealt{lon66,sch68,croom04,rich06b}). The period between $z\\sim2$ and $3$ marks the peak of quasar activity in the Universe. Understanding the processes which caused this ramping up of activity in the early Universe and what is responsible for its reversal is a major goal in QSO science. The fact that the \\civ\\ line is visible in optical spectra over this entire range potentially makes it an attractive probe of QSOs at these epochs.  In section~\\ref{sec:vir_mass} we review virial SMBH mass estimation and the use of the \\civ\\ line with this technique. Section~\\ref{sec:data} gives a brief description of the three datasets used in this work. In section~\\ref{sec:results} we present the results of our fitting and make a comparison with a previous analysis of the \\mgii\\ line in the same sample of quasars from \\citet{me2}. Motivated by the results in \\citet{me2}, section~\\ref{sec:disp} describes an investigation of the dispersion in \\civ\\ line widths as a function of redshift and luminosity. In sections~\\ref{sec:vir_wrap} and~\\ref{sec:geometry} we discuss our results with respect to virial SMBH mass estimation and BLR geometry. Details of the analytic procedure, including a discussion of our line fitting analysis and a prescription for removing broad-absorption lines (BALs) from our data, can be found in appendices~\\ref{sec:fitting} and~\\ref{sec:bal} respectively. Throughout this paper we assume a flat $(\\Omega_{\\rm m},\\Omega_{\\Lambda})=(0.3,0.7)$, $H_{0}=70\\,{\\rm km\\,s}^{-1}\\,{\\rm Mpc}^{-1}$ cosmology. ",
        "conclusions": "We have performed an extensive analysis of the \\civ\\ line width distribution in QSO spectra from the 2SLAQ, 2QZ and SDSS. We reviewed the three most common methods for fitting \\civ\\ emission in QSO spectra and performed a detailed comparison between these, using both composite spectra and individual fits to our whole dataset. Based on these results we have chosen the procedure that we believe is the most physically motivated and least biased to employ for our analysis. Furthermore we have developed a procedure for identifying absorption features in spectra and removing BAL systems from the dataset.  Applying our routine to spectra from the SDSS, 2QZ and 2SLAQ surveys we have measured the \\civ\\ line width distribution for QSOs in the redshift range $1.5 < z < 3.3$ spanning a magnitude range of $-24 > M_r(z=2.5) > -29$. We find that the line width vs. luminosity plot (Fig.~\\ref{fig_M_lw_c}) shows many similarities with the equivalent plot for \\mgii. However, this appears to be where the similarity ends.  We compare the \\civ\\ and \\mgii\\ line widths calculated in spectra that have both lines, and find considerable scatter in the comparison. We find that the dispersion in \\civ\\ line width is smaller than that for \\mgii\\ and does not show the same trend with luminosity that \\mgii\\ exhibits.  These results are discussed in terms of virial SMBH mass estimation. We show that, for both \\mgii\\ and \\civ\\ estimators, the line width term contributes considerably less dynamic range to the resulting mass estimate than the quoted error on the estimate. The dynamic range found in virial SMBH mass estimates comes (almost entirely) from their luminosity term, and this is solely responsible for the consistency of virial masses based on the two lines since the line widths do not correlate.  Finally the results are discussed in terms of BLR dynamics. We show that given the small scatter in \\civ\\ line widths the \\civ\\ BLR cannot be a flat disk. We parametrise models for a hybrid BLRs which include a random/isotropic velocity component and show how our results can be used to constrain these models."
    },
    "1005/1005.2658_arXiv.txt": {
        "abstract": "{ We have investigated some of the properties of dense sub-nuclear matter at the crustal region (both the outer crust and the inner crust region) of a magnetar. The relativistic version of Thomas-Fermi (TF) model is used in presence of strong quantizing magnetic field for the outer crust matter. The compressed matter in the outer crust, which is a crystal of metallic iron, is replaced by a regular array of spherically symmetric Wigner-Seitz (WS) cells. In the inner crust region, a mixture of iron and heavier neutron rich nuclei along with electrons and free neutrons has been considered. Conventional Harrison-Wheeler (HW) and Bethe-Baym-Pethick (BBP) equation of states are used for the nuclear mass formula. A lot of significant changes in the characteristic properties of dense crustal matter, both at the outer crust and the inner crust, have been observed. \\PACS{ {97.60.Jd}{ Neutron stars} \\and {71.70}{Landau levels} \\and {26.60.GJ}{Neutron star crust} \\and {26.69.Kp}{Equation of State of neutron-star matter} } } ",
        "introduction": "\\label{intro} Magnetars are the most exotic stellar objects, believed to be strongly magnetized young neutron stars. The surface magnetic field for such objects are observed to be $\\geq 10^{15}$G \\cite{R1,R2,R3,R4}. Then it is quite likely that the field at the interior, even at the inner crust region may be stronger than the surface value (can be predicted theoretically by scalar Virial theorem). If the internal field strength is happened to be so high, then most of the physical and chemical properties of dense stellar matter of the magnetars must change significantly from the conventional picture \\cite{R5,R6,R7}. A lot of investigations have already been done on the effect of strong quantizing magnetic field on various physical properties of dense stellar matter inside neutron stars as well as quark matter inside quark stars or hybrid stars, including the effect on quark-hadron phase transition at the core region of a compact neutron star \\cite{R8}. The effect of such strong magnetic field on various elementary processes inside neutron stars and quark stars or hybrid stars have also been studied. The effect of strong quantizing magnetic field on the $\\beta$-equilibration among the constituents have also been investigated \\cite{R9}. It has also been shown that strong quantizing magnetic field acts like a catalyst to generate fermion mass dynamically, i.e., chiral  symmetry breaking occurs in presence of strong quantizing magnetic field \\cite{R10,R11,R12}.  In this article we present our investigation on the effect of strong magnetic field on the crustal matter of magnetars. The work is divided into two parts: in the first part, based on one of our very recent work \\cite{R13}, we have investigated the effect of strong quantizing magnetic field on the outer crust matter. In the second part, we have studied the properties of compact sub-nuclear matter at the inner crust region in presence of such strong quantizing magnetic field.  The paper is organized in the following manner: In section 2, the effect of strong quantizing magnetic field on the crustal matter of magnetars is discussed. In the outer crust region, the matter with dense crystalline structure of metallic iron at sub-nuclear density are replaced by an array of spherically symmetric WS cells with positively charged nuclei at the centre surrounded by non-uniform dense electron gas with over all charge neutrality. In the inner crust we have used the conventional HW and BBP equation of states for the nuclear matter \\cite{RR14,R14,R15}, consisting of iron along with some more heavier neutron rich nuclei. In one of our forth coming article we shall report the effect strong magnetic field on inner crust matter using some modern type nuclear mass formulas and make a comparative study with those models. Here we have presented with detailed numerical computation the effect of strong magnetic field on the inner crust matter, which as mentioned above is assumed to be a mixture of iron, some heavier neutron rich nuclei, which we found to be specially true in presence of strong quantizing magnetic field, electrons and free neutrons. The presence of free neutrons are considered beyond neutron drip density. Finally, in the last section, we have given the conclusions and discussed the importance and future prospects of the present work. ",
        "conclusions": "\\label{sec:4} In the first part of this article we have presented our investigation on the properties of dense outer crust matter of the magnetars. We have replaced the dense metallic iron crystal by a regular array of spherically symmetric WS cells. It has been observed that the radius of each cell decreases with the magnetic field strength. In this article, however, we have not considered the magnetic field induced deformation of WS cells, which is quite likely if the magnetic field is strong enough. In future we shall present this important issue with a completely different approach.  It has also been noticed that the upper limit of Landau quantum number is a function of positional coordinate of the electron with which it is associated within the WS cells. We have observed that at the surface region, for all the values of magnetic field strength, this upper limit becomes identically zero. Which actually means that the electrons near the WS cell surface are strongly polarized in the opposite direction of external magnetic field. Whereas, for $B > 10^{15}$G, they are polarized at every points within the cells.  It has been observed that the electron density within the cells increases with the increase in magnetic field strength. Further, for all the values of magnetic field strength, the electron density is maximum near the nuclear surface ($r=r_n$) and minimum at the WS cell boundary ($r=r_s$).  The electron kinetic pressure is found to be positive near the central portion of WS cell but is negative near the surface region. There is a position within the cell at which the kinetic pressure vanishes. The position of this point changes with the strength of magnetic field. We have also studied the variations of kinetic energy, electron-nucleus interaction energy, electron-electron direct potential energy and electron-electron exchange interaction part within the cells. We have shown the variation of these quantities within the WS cells for a number of magnetic field strengths.  In the second part of this work, we have investigated some of the properties of inner crust matter of magnetars. We have used HW and BBP equation of states for the nuclear mass formula. It has been observed from both the models, that for a stable inner crust matter, the nuclei present must be heavier than iron and much more neutron rich. The heaviness is more in the case of BBP equation of state. We have also noticed that for low and moderate values of magnetic field strength, the variation of mass number and the corresponding atomic number with magnetic field is not so significant. Whereas, for $B \\geq 10^{15}$G, when electrons occupy only the zeroth Landau level, then much more heavier neutron rich nuclei are formed in the inner crust region. It is found that high magnetic field behaves like a catalyst to generates heavy neutron rich nuclei.  We have observed that initially the electron density increases with the increase in mass number, but as soon as free neutrons appear in the system, the electron density decreases and saturates to some constant value which depends on the magnetic field strength. In the case of HW equation of state, free neutron density does not depend on the strength of magnetic field, whereas, for BBP case, because of chemical equilibrium condition, the free neutron density depends very weakly on the magnetic field strength. We have noticed that in the case of BBP equation of state the overall qualitative difference is because of chemical equilibrium among the constituents.  The total baryon density rises sharply like an avalanche for the value of $A$ at which free neutrons appear in the system. However, for BBP equation of state, because of chemical equilibrium condition, the rise is not so sharply visible for a given magnetic field $B$.  The qualitative nature of equation of states are almost identical. It is found that in presence of strong magnetic field, the inner crust matter becomes mechanically stable (with the positive value of kinetic pressure) only at very high density.  We therefore believe that to study various properties of dense matter associated with magnetars, one must consider all these significant changes from the conventional scenario. Further, we expect that the strong magnetic field, if present well within the inner crust of magnetars, must affect the super-fluidity of cold neutron matter."
    },
    "1005/1005.3067.txt": {
        "abstract": "% context heading (optional) % {} leave it empty if necessary {It was proposed earlier that the  relativistic ejections observed in microquasars could be produced by violent magnetic reconnection episodes at the inner disk coronal region (de Gouveia Dal Pino \\& Lazarian 2005). } % aims heading (mandatory) %{} {Here we revisit this model, which employs a standard accretion disk description and fast magnetic reconnection theory, and discuss the role of magnetic reconnection and associated heating and particle acceleration in different jet/disk accretion systems, namely young stellar objects (YSOs), microquasars, and active galactic nuclei (AGNs).} % methods heading (mandatory) %{} {In microquasars and AGNs, violent reconnection episodes between the magnetic field lines of the inner disk region and those that are anchored in the black hole are able to heat the coronal/disk gas and accelerate the plasma to relativistic velocities through a diffusive first-order Fermi-like process within the reconnection site that will produce intermittent relativistic ejections or plasmons.} % results heading (mandatory) %{} {The resulting power-law electron distribution is compatible with the synchrotron radio spectrum observed during the outbursts of these sources. A diagram of the magnetic energy rate released by violent reconnection as a function of the black hole (BH) mass spanning $10^9$ orders of magnitude shows that the magnetic reconnection power is more than sufficient to explain the observed radio luminosities of the outbursts from microquasars to low luminous AGNs. In addition, the magnetic reconnection events cause the heating of the coronal gas, which can be conducted back to the disk to enhance its thermal soft x-ray emission as observed during outbursts in microquasars.  The decay of the hard x-ray emission right after a radio flare could also be explained in this model due to the escape of relativistic electrons with the evolving jet outburst. In the case of YSOs a similar magnetic configuration can be reached that could possibly produce observed x-ray flares in some sources and provide the heating at the jet launching base, but only if violent magnetic reconnection events occur with episodic, very short-duration accretion rates which are $\\sim 100-1000 $ times larger than the typical average accretion rates expected for more evolved (TTauri) YSOs. %This study is  a first step and  multi-dimensional %numerical models will be necessary to confirm the present results. } % conclusions heading (optional), leave it empty if necessary %{} {} % ",
        "introduction": "Supersonic jets are observed in several astrophysical systems, such as active galactic nuclei (AGNs), neutron star and black hole X-ray binaries, and low-mass young stellar objects (YSOs), and are probably also associated with gamma-ray bursts.  The study of their origin, structure, and evolution helps to shed light on the nature of their compact progenitors, as they carry angular momentum, mass, energy, and magnetic field away from the sources.   The currently most accepted paradigm for jet production is based on the magneto-centrifugal acceleration out of a magnetized accretion disk that surrounds the central source. Firstly proposed by  Blandford \\& Payne (1982; see also Lynden-Bell 1969; Blandford \\& Rees 1974; and Lovelace 1976, where these ideas initially germinated), this basic scenario for jet launching has been object of extensive analytical and numerical investigation (see e.g.,  McKinney \\& Blandford 2009; Shibata 2005; de Gouveia Dal Pino 2005 for reviews). However, though considerable progress has been achieved in the comprehension of the possible origin of the magnetic fields that must permeate the accretion disk and the mechanism of angular momentum transport that allows the accretion to occur through magnetorotational turbulence (Balbus \\& Hawley 1998), questions related to jet stability, the nature of the coupling between the central source magnetosphere and the disk field lines, and the quasi-periodic ejections that are often associated to these jets, are  still debated.  Relativistic jets from stellar-mass black holes of binary stars emitting X-rays, also called microquasar jets (or BHXRB jets), are scaled-down versions of AGN (or quasar) jets, typically extending for $\\sim$ 1pc and probably powered by the accreting, spinning black hole. Despite the enormous difference in scale, both classes share several similarities in their physical properties. However, because the characteristic times of the matter flow are proportional to the black hole mass, the accretion-ejection phenomena in microquasars end sooner and are about $10^{-7} - 10^{-5}$ faster than analogous phenomena in quasars. This fact and their closer proximity to us (they are generally observable within the Galaxy) make the microquasars easier to investigate compared to the AGNs (Mirabel \\& Rodrigues 1999).  Although individual systems are complex and peculiar when looked at in detail, there are common features to all classes of BHXRBs (e.g, Remillard \\& McClintock, 2006). According to their x-ray emission (2-20 keV), they show basically two major states: a quiescent and an outburst state. The former is characterized by low x-ray luminosities and hard non-thermal spectra. Usually, transient BHXRBs exhibit this state for long periods, which allows one to obtain the physical parameters of the system as the spectrum of the secondary star becomes prominent.  On the other hand, the outburst state corresponds to intense activity and emission and can be sub-classified in three main active and many intermediary states. According to Remillard \\& McClintock (2006), the three main active states are the thermal state (TS), the hard state (HS) and the steep power law state (SPLS). These states are usually explained as changes in the structure of the accretion flow. During the TS, for example, the soft x-ray thermal emission comes from the inner region of the thin accretion disk that extends until the last stable orbits around the black hole. On the other hand, during the HS the observed weak thermal component suggests that the disk is truncated at a few hundreds/thousands gravitational radii. The hard x-ray emission measured during this state is often attributed to inverse Compton scattering of soft photons from the outer disk by relativistic electrons in the hot inner region of the system (e.g. Remillard \\& McClintock, 2006; Malzac, 2007).  A widely observed example from radio to x-rays is the microquasar GRS 1915+105. Located at a distance of $\\sim 12.5$ kpc and probably with a 10--18 solar mass black hole in the center of the binary system, it was the first galactic object to show evidence of superluminal radio ejection (Mirabel \\& Rodrigues, 1994; Mirabel \\& Rodrigues, 1998). Dhawan et al. (2000) have distinguished two main radio states of this system, a plateau and a flare state. During the plateau state the RXTE (2--12 keV) soft X-ray emission is weak, while the BATSE (20--100 keV) hard X-rays are strong and the radio flat spectrum is produced by a small scale nuclear jet. On the other hand optically thin ejecta are superluminally expelled up to thousands of AU during the flare phase and the radio spectral index is between 0.5 and 0.8. The soft X-rays also flare during this phase and exhibit a high variability, while the hard X-rays fade for a few days before recovering. In terms of the x-ray spectral states, several works have verified that the radio flares of this source occur during the SPLS (e.g., Fender et al., 2004) and the x-ray emission of both the plateau and the flare state are different manifestations of the SPLS (e.g., Reig et al., 2003).  The examination  of  radio and x-ray observations for other microquasars (e.g., XTE J1859+226, XTE J1550-564) also suggests that the ejection of relativistic matter happens when the source is very active and in the SPLS (e.g., Hannikeainen et. al, 2001; Brocksopp et. al, 2002; Fender et al., 2004; McClintock et. al, 2007). According to Fender et al. (2004), the origin of the optically thin emission in radio is due to shock waves that are formed in the jet when the system passes from a `hard' SPLS to a `soft' SPLS. However, what generates these shock waves or the triggering mechanism for the ejections of matter is unclear.  In 2005, de Gouveia Dal Pino \\& Lazarian (see also de Gouveia Dal Pino 2006) proposed a mechanism that can be responsible for the initial acceleration of the plasma jet to relativistic speeds in the case of  the microquasar GRS1915+105. Their scenario is related to violent reconnection episodes between the magnetic field lines of the inner disk region and those that are anchored in the black hole. We here revisit this model and argue that it could be responsible for the transition from the `hard' SPLS to the `soft' SPLS seen in other microquasars. We also extend it to the AGN jets and briefly discuss its role for thermal YSO jets (see also de Gouveia Dal Pino 2006 and de Gouveia Dal Pino et al. 2009 for earlier discussions).   We note that an interesting correlation between the radio and the (hard) X-ray luminosity has been reported by Laor \\& Behar (2008) for sources spanning $10^{10}$ orders of magnitude in mass (from magnetically active stars to some galactic black holes and radio quiet AGNs), and recently extended to dwarf nova outbursts by Soker \\& Vrtilek (2009). This correlation holds only for sources for which $L_R/L_X < 10^{-5}$ and suggests that a common mechanism may be responsible for  the radio emission in these highly different objects. Because in magnetically active stars the emission is related to the coronae, these authors have argued that the same might be occurring in some BHXRBs and radio quiet AGNs, i.e.,  the radio emission could also come from magnetic activity in the coronae above the accretion disk, just like in the scenario that we investigate here in detail (and which was previously suggested in de Gouveia Dal Pino 2006, and de Gouveia Dal Pino et al. 2010a, 2010b).  Earlier studies have also revealed a correlation between the radio and the X-ray luminosity of BHXRBs in their quiescent \"low/hard\" state and low-luminosity AGNs with nuclear radio emission or weak radio jets  (Gallo et al. 2003; Merloni et al. 2003; Falcke et al. 2004). \\footnote{We notice that these correlations are often interpreted in the context of (advection dominated accretion flow) ADAF models  where a  geometrically thick and optically thin accretion disk  with small mass flow rate is invoked to explain the low/hard X-ray emission accompanied by radio jet activity. In these models a geometrically thin disk structure in the inner region would be recovered when the system switches to a high-accretion flow rate, hence producing  a high/soft emission with suppression of the jet activity. In the present model we examine these objects mostly during their outburst phases rather than during their more quiescent states. Besides, because the reconnection mechanism is operating at the interface region between the inner disk and the central source, the choice of a geometrically thin accretion disk (rather than a  thick one) is not a crucial point for our model.}    As also remarked by Soker \\& Vrtilek, suggestions for the presence of coronae above accretion disks are not new (e.g., Liang \\& Price 1977, Liang 1979, Galeev et al. 1979; Done \\& Osborne 1997; Wheatley \\& Mauche 2005; Cao 2009), and the connection between coronae and jets were previously proposed as well (e.g., Fender et al. 1999; Markoff et al. 2005; de Gouveia Dal Pino \\& Lazarian 2000, 2001, 2005). If the magnetic activity in erupting accreting disks is similar to that in stars, then coronal mass ejections,  as in the Sun, could be expected. Hence magnetic flares similar to those in active stars might be a potential mechanism operating at the jet launching region in a variety of systems, from young stellar objects to BHs (de Gouveia Dal Pino 2006, de Gouveia Dal Pino et al. 2009; Soker 2007;  Soker \\& Vrtilek 2009).   %The %idea that jet outbursts may be launched by magnetic fields reconnection % events similar to those in the sun is not new, e.g., de Gouveia Dal Pino \\& %Lazarian (2005; de Gouveia Dal Pino 2006; de Gouveia Dal Pino et al. 2008).      This paper is organized as follows: in Sect. 2 we briefly describe de Gouveia Dal Pino \\& Lazarian's model. In Sect. 3 we discuss it in the context of the microquasars and in Sect. 4 we propose its generalization to other classes of relativistic jets, extending it to low-luminosity AGN jets. In Sect. 5 we briefly discuss the reconnection model in the context of the YSO jets and finally, in Sect. 6, we draw our conclusions. ",
        "conclusions": "We revisited the model proposed by de Gouveia Dal Pino \\& Lazarian (2005) for the production of relativistic ejections of microquasar GRS 1915+195. Within this scenario, the initial acceleration of the jet plasma to relativistic speeds is related to violent reconnection episodes between the magnetic field lines of the inner disk region and those that are anchored in the black hole. This might happen whenever a large scale magnetic field (of opposite polarity to that anchored in the BH) arises from the inner accretion disk (probably amplified by dynamo action) so that the ratio between the effective disk pressure and the magnetic pressure ($\\beta$) decreases to values near 1 or below and the accretion rate ($\\dot{M}$) approaches the Eddington rate. In Sect. 3 we extended this model to other microquasars and argued that it could be responsible for the transition from the `hard' SPLS to the `soft' SPLS seen in other microquasars.  In this scenario, the high accretion rate right before a radio flare results in a soft x-ray emission with a luminosity around $\\sim 10^{39}$ erg/s (Sect. 3.4). The magnetic power released by violent magnetic reconnection events is mainly used to heat the coronal gas to temperatures of $\\sim 10^9$ K (see Eq. 8) and heat the superficial disk gas by thermal conduction along the field lines causing  enhancement and variability of the soft x-ray flux (section 3.5). A substantial fraction of $\\dot{W}_B$ is also expected to be used to accelerate particles that produce the relativistic blobs that emit radio-synchrotron-radiation with spectral index of $0.75$ if the electrons are first accelerated in the reconnection zone by  first-order Fermi like process (Sect. 3.3; de Gouveia Dal Pino \\& Lazarian 2005). Further out the relativistic fluid may be also re-accelerated behind shocks, which are expected to be formed by the magnetic plasmons that erupt from the reconnection zone. Behind these shocks a standard first-order Fermi acceleration must also produce  a synchrotron spectrum $S_{\\nu} \\propto \\nu^{-0.5}$. Both radio spectral indices are consistent with the observed spectral range during the flares ($\\alpha_{R} \\simeq -0.2$ to $-1$, Dhawan et al., 2000; Hannikainen et al., 2001). The hard x-ray emission is produced by inverse Compton scattering of the soft x-ray emission by the accelerated relativistic electrons and right after a radio flare, the hard x-ray flux should decrease, because most of these electrons are accelerated away from the disk (Sect. 3.4).  In Sect. 4 we generalized this model to AGNs. We built a diagram that compares  the calculated power released during violent magnetic reconnection versus the sources black hole masses with the observed outburst radio emission from microquasars and AGNs (Fig. 3). The results indicate  that the coronal  magnetic activity  near the jet launching region can explain the observed emission of relativistic radio blobs from microquasars to low-luminous AGNs (LLAGNs), spanning over $10^9$ orders of magnitude in  mass of the sources. This correlation does not hold for radio-loud AGNS. These  would require reconnection events with super-Eddington accretion to explain the formation and  emission of relativistic blobs as due to reconnection. This is possibly because their surroundings are much denser and then \"mask\" the emission due to  coronal magnetic activity at sub-Eddington rates. In this case, particle re-acceleration behind shocks will probably prevail further out in the jet launching region and will be responsible for the radio emission.  %This super-Eddington regime would manifest itself only in a short time interval %until the partial destruction of the magnetic flux by MR would relax this extreme %configuration. This will be investigated in a future paper.  %Finally, we notice that the mechanism investigated here for relativistic blobs production is compatible with the proposed unified scenario for astrophysical jet production based on the magneto-centrifugal scenario (e.g., Blandford \\& Payne 1982).   A similar magnetic reconnection model applied to the jet launching region of YSO  jets suggests that  the observed X-ray flares in the sample of COUP sources (Favata et al. 2005) could be explained by violent magnetic reconnection events in the magnetosphere-inner disk coronal region only if high episodic accretion rates are achieved with $dM/dt  \\simeq 100-1000 <dM/dt>$ and  the inner disk radius approaches the stellar radius during these events. These high accretion rates, however, would last only very briefly, i.e.,  not for times much longer than the typical reconnection time (Eq. 13).  We also note that the predicted range of values for the  magnetic reconnection power of YSOs naturally lies in the left inferior part of the magnetic reconnection power versus source mass diagram of Fig. 3. This  is consistent  with the interpretation that the emission processes investigated above in all these classes of sources would  be  mainly associated with magnetic activity in the coronae and therefore would be nearly independent of the intrinsic physics of the central source and the accretion disk.  Finally, we remark that the simple analytical study carried out here has allowed us to derive  estimates of the maximum amount of energy that can be released from violent magnetic reconnection events in the inner region of several accretion disk-coronal-jet systems, as well as its potential effects on these systems, but this is only a first step, and these results must be tested with more realistic multi-dimensional numerical models which are able to capture the intrinsic non-linearity and non-steadiness character of the problem."
    },
    "1005/1005.4248_arXiv.txt": {
        "abstract": "{\\it Spitzer} and {\\it AKARI} observations have found that polycyclic aromatic hydrocarbons (PAHs) are present in nearby elliptical galaxies, but their spatial distributions are still unknown. In order to investigate their distributions, we performed deep spectral mapping observations of the PAH-detected elliptical galaxy NGC~4589, a merger remnant with a minor-axis optical dust lane. As a result, we obtain clear evidence that the PAH 11.3 $\\mu$m emission comes predominantly from the dust lane of the galaxy. We also detect molecular hydrogen line emissions from the dust lane. The PAH 17 $\\mu$m emission is distributed differently from the PAH 11.3 $\\mu$m emission, and more similarly to the dust continuum emission. From their distinctive distributions, we suggest that the PAHs responsible for the 11.3 $\\mu$m feature are secondary products through the evolution of the ISM brought in by the merger. ",
        "introduction": "Many elliptical galaxies are not ISM-free, which contain a considerable amount of cool dust (e.g. Knapp et al. 1989; Goudfrooij \\& de Jong 1995; Temi et al. 2004; Temi et al. 2007). Early {\\it Spitzer} IRS observations have found that even polycyclic aromatic hydrocarbons (PAHs) are present in elliptical galaxies (Kaneda et al. 2005). The spectra show that the PAH 11.3 $\\mu$m feature is notably strong relative to usually the strongest 7.7 $\\mu$m feature, which might be typical in evolved systems such as elliptical galaxies, i.e. the dominance of neutral PAHs over ionized ones due to very soft radiation fields from evolved stars. In addition to the PAH features, a series of H$_{2}$ rotational lines as well as ionic fine-structure lines (e.g. [NeII] and [NeIII]) are detected from a large fraction of the elliptical galaxies, implying the presence of warm molecular and ionized gases \\citep{Kan08a}. Their presence seems to be inconsistent with the dominance of neutral PAHs irradiated by soft radiation field.  What are the origin and fate of the PAHs in elliptical galaxies? How are they related with the warm molecular and ionized gases as well as the cool dust? How do their properties change from the center to outer regions? Spatial information is indispensable to disentangle various competing components and tackle these issues. To pursue follow-on studies of the early {\\it Spitzer} discoveries, we performed deep spectral mapping observations of elliptical galaxies with the {\\it Spitzer} IRS.  In this Letter, we report the result of NGC~4589, a nearby E2 elliptical galaxy with a minor-axis optical dust lane \\citep{Mol89}. NGC~4589 has a complex stellar rotation field, although its morphology shows the smooth optical profile following the de Vaucouleurs law \\citep{Mol89}, which supports that the galaxy is a relatively old merger remnant. Recently, a core-collapse supernova (SN Ib 2005cz) was observed in this genuine elliptical galaxy \\citep{Hak08}; SN Ib 2005cz is likely to be the end product of one of young stars, which were produced by the galaxy merger about $10^8$ years ago \\citep{Kaw10}. NGC~4589 is an X-ray-emitting elliptical galaxy with a low-ionization nuclear emission-line region (LINER) nucleus \\citep{OSu01, Dre85, Wil85}.  Our earlier IRS observation detected the PAH features from the galaxy in a staring mode but not in a mapping mode, and therefore we could not spatially resolve them \\citep{Kan08a}. After the {\\it Spitzer} detection of the PAH features, their spatial distributions were estimated from the {\\it AKARI} 11 $\\mu$m and 15 $\\mu$m band images, while that of the far-infrared (IR) dust was obtained from the {\\it AKARI} 90 $\\mu$m image \\citep{Kan08b}. They all show significant deviations from the smooth stellar distribution, from which Kaneda et al. (2008b) suggested that the PAHs and far-IR dust are mostly of an interstellar origin, rather than of a stellar origin. In this Letter, from the {\\it Spitzer}/IRS spectral mapping of NGC~4589, we clearly find that the PAH 11.3 $\\mu$m emission comes from the minor-axis dust lane, while the PAH 17 $\\mu$m emission is distributed differently from the PAH 11.3 $\\mu$m emission, and more similarly to the far-IR dust emission. ",
        "conclusions": "With the {\\it Spitzer}/IRS, we performed deep spectral mapping observations of NGC~4589, a merger remnant with a minor-axis optical dust lane. As a result, we obtain clear evidence that the PAH 11.3 $\\mu$m emission comes predominantly from the optical dust lane of the galaxy. We also detect molecular hydrogen line emissions from the dust lane. The distribution of the PAH 17 $\\mu$m emission is extended more widely in directions similar to that of the far-IR dust continuum emission. From their distinctive distributions, we suggest that the PAHs emitting the 11.3 $\\mu$m feature are secondary products through the evolution of the ISM brought in by the merger, while the distributed dust and large PAHs emitting the 17 $\\mu$m feature are relics of the merger."
    },
    "1005/1005.3098.txt": {
        "abstract": "{Optical spectra of the peculiar supergiant 3\\,Pup taken in 1997--2008 are used to study the spectral peculiarities and velocity field in its atmosphere. The profiles of strong FeII lines and of the lines of other ions have a specific shape: the wings are raised by emissions, whereas the core is sharpened by a depression. The latter feature becomes more pronounced with the increasing line strength, and the increasing wavelength. Line profiles are variable: the magnitude and sign of the absorption asymmetry, and the blue--to--red emission intensity ratios vary from one spectrum to another. The temporal Vr variations are minimal for the forbidden emissions and sharp shell cores of the absorption features of FeII(42), and other strong lines of iron-group ions. Therefore the average velocity for the above lines can be adopted as the systemic velocity: V$_{\\rm sys} = 28.5 \\pm 0.5$\\,km/s. The weakest photospheric absorptions and photospheric MgII, SiII absorptions exhibit well-defined day-to-day velocity variations of up to 7\\,km/s. Quantitative spectral classification yields the spectral type of A2.7\\,$\\pm$\\,0.3~Ib. The equivalent widths and profiles of H$\\delta$ and H$\\gamma$, and the equivalent width of the OI~7774\\,\\AA{} triplet yield an absolute magnitude estimate of M$_{\\rm v}=-5\\lefteqn{.}^m5 \\pm 0\\lefteqn{.}^m3 $, implying the heliocentric distance of 0.7\\,kpc. }  \\authorrunning{Chentsov, Klochkova  \\&  Miroshnichenko} \\titlerunning{Spectral variability of 3\\,Pup} ",
        "introduction": "The nature of the bright star 3\\,Pup (HD\\,62623\\,=\\,HR\\,2996\\,=\\,HIP\\,37677) has long remained unclear. In the Bright Star Catalog [1] this star is listed as a A2\\,Iabe--supergiant. However, 3\\,Pup was also viewed as an object related to stars evolving towards planetary nebulae with binary core~[2]. The spectroscopic study of Plets~et~al.~[3], who found the star's atmosphere to have a close-to-solar chemical composition, provided new evidence for the massive supergiant nature of 3\\,Pup. However, the star has a property that is uncharacteristic to massive A-type supergiants: a circumstellar dust shell. The shell shows up primarily in the form of IR--excess (3\\,Pup was identified with the IR--source IRAS\\,07418$-$2850) and specific two-peaked emission features in the optical spectrum.  The star's magnitudes are B\\,=\\,4$\\lefteqn{.}^{\\rm m}12$ and V\\,=\\,3$\\lefteqn{.}^{\\rm m}98$, and the color excess E(B--V)\\,=\\,0$\\lefteqn{.}^{\\rm m}08$~[4]. Its Galactic and equatorial J2000 coordinates are l\\,=\\,244$\\lefteqn{.}{\\degr }4$, b\\,=\\,--2$\\lefteqn{.}{\\degr}5$ and $\\alpha$\\,=\\,$07^{\\rm h} 43^{\\rm m} 48^{\\rm s}$, $\\delta =-28{\\degr}$57${'}$\\,17${''}$, respectively. The recently refined parallax is $\\pi$\\,=\\,0.59$\\pm$0.17\\,mas~[5]. Because of its high brightness 3\\,Pup was chosen as one of the several hundred stars used as the basis of Harvard spectral classification. However, whereas Cannon~[6] assigned to it the same spectral type A2\\,pec as to $\\alpha$\\,Cyg (the peculiarity of both stars showed up in the form of too narrow and weak hydrogen lines, i.e., both stars are supergiants), Maury~[7] noticed that 3\\,Pup has a somewhat lower temperature and luminosity than $\\alpha$\\,Cyg, A2--A3 and Ib--II, respectively, in the currently adopted notation.  Merrill~[8] found  some emission lines in the spectrum of 3\\,Pup: a double H$\\alpha$ line with a relatively stronger red component (a P Cyg III type profile according to Beals~[9]), [OI]\\,1F\\,6300, 6364\\,\\AA{}, FeII(40)\\,6433, 6516\\,\\AA. He also suspected the presence of emission features in some other FeII lines. Emission features in the profiles of H$\\beta$, H$\\gamma$~[10] and NaI(1)~doublet~[3] were noticed only 60 years later. The known infrared emission features include the strong CaII\\,(2)\\,8498, 8542, 8562\\,\\AA{} triplet~[11], and the [CaII]\\,1F 7291, 7324\\,\\AA{} doublet~[12] (these lines were erroneously attributed to [OII]). As for absorption lines, the eye estimates of their intensities in the 3120--6565\\,\\AA{} wavelength interval~[13, 14] confirmed a lower excitation temperature of 3\\,Pup, compared to that of $\\alpha$Cyg, as it followed from classification of Maury, and revealed shell signatures in the star's spectrum. In particular, the MgII\\,(4) and SiII\\,(1,2,3) absorptions in the spectrum of 3\\,Pup are weaker than the same absorptions in the spectrum of $\\alpha$\\,Cyg, whereas the FeII\\,(42) lines are enhanced compared to FeII\\,(37, 38, 49) lines. The differences between the spectra of a supergiant and a shell star are readily seen in Plate~22 in the atlas of Morgan et al.~[15], where they are illustrated on the examples of $\\alpha$\\,Cyg~(A2~Ia) and 17\\,Lep (Ap sh). The temperature of 3\\,Pup is close to the temperatures of both stars, but its spectrum is closer to that of 17\\,Lep by a number of features (the depth of TiII and FeII lines and the widths of hydrogen lines).  It became clear from the first inspection of the available spectra of 3\\,Pup that it must be a very promising object for high-resolution spectroscopy: the star is bright and its spectrum abounds in various, mostly weak, but well-defined features. But however strange it may appear, its optical and near infrared spectrum has not yet been thoroughly described. This circumstance led us to include 3\\,Pup into our list of program stars with circumstellar shells for a detailed spectroscopic analysis in order to study the expected spectral variability. ",
        "conclusions": "In the case of FeII lines in the spectrum of 3\\,Pup there is an evident contribution of the shell, which gives the profiles their specific shapes: the wings are raised by emissions, whereas the core is sharpened. The latter is apparent even for the absorptions with depths R\\,$\\approx$\\,20 at $\\lambda>4300$\\AA{} and becomes more pronounced with increasing line strength. Sharpened cores are typical of Balmer lines except for H$\\alpha$.  The profiles of the CaII\\,(2) triplet have a ``twostep'' shape: a narrow double peak is appreciably redshifted with respect to the relatively broader ``pedestal''. The radial velocity of the peak is close to the velocities of forbidden emission lines, whereas that of the lower component of the profile is close to the velocities determined from the permitted emissions.  Line profiles vary with time: both the magnitude and the sign of the absorption asymmetry, and the blue--to--red emission intensity ratios vary from one spectrum to another. The profiles of all forbidden emissions in the recorded region of the spectrum have the same shapes and widths and yield the same radial velocity within the errors.  Radial velocity variations of forbidden emissions and sharp envelope cores of FeII\\,(42) absorptions and other strong iron-group ion lines are minimal and the velocities of these lines are close to each other. The mean velocity averaged over the above lines can be adopted as the radial velocity of the system as a whole: V$_{\\rm sys}$\\,=\\,28.5$\\pm$0.5\\,km/s.  The weakest photospheric absorptions (with residual intensities r$\\rightarrow$100) and the photospheric absorptions MgII, SiII exhibit bona fide day-to-day velocity variations amounting to 7\\,km/s.  Shallow FeI, FeII, and TiII absorptions and MgI, MgII lines in the 4280--4700\\,\\AA{} wavelength interval yield a spectral type of A2.7$\\pm$0.3~Ib. The equivalent widths and profiles of H$\\delta$ and H$\\gamma$, and the equivalent width of the OI\\,7774\\,\\AA{} triplet yield an absolute magmnitude of Mv\\,=\\,$-5.5\\pm  0.3^{\\rm m}$.    \\newpage"
    },
    "1005/1005.4886_arXiv.txt": {
        "abstract": "We present {\\it Spitzer Space Telescope} observations of the ``peculiar variable'' DZ~Cru, identified by \\citeauthor[][]{rushton} (2008, MNRAS, 386, 289) as a classical nova. A dust shell, on which are superimposed a number of features, is prominent in the 5--35\\mic\\ range some 4~years after eruption. We suggest that the dust in DZ~Cru is primarily hydrogenated amorphous carbon in which aliphatic bands currently predominate, and which may either become predominantly aromatic as the dust is photo-processed by ultraviolet radiation from the stellar remnant, or more likely completely destroyed. ",
        "introduction": "DZ~Cru (also known as Nova Cru 2003 and the ``peculiar variable'' DZ Cru) was discovered in 2003 August \\citep{tabur}. It was initially reported as a nova but spectroscopic observations gave rise to doubts about its nova status, and comparisons were drawn with the unusual variable V838~Mon \\citep{dellavalle}. \\cite{rushton} presented  1--2.5\\mic\\ infrared (IR) spectra obtained $\\sim1-1.5$~years after discovery. They found \\pion{H}{i}, \\pion{O}{i} and \\fion{N}{i} emission lines superimposed on a hot dust continuum and concluded that DZ~Cru was indeed a classical nova at distance $\\sim9$~kpc (although this is subject to considerable uncertainty).  Here we present observations of DZ~Cru with the {\\it Spitzer} Space Telescope \\citep{spitzer,gehrz-spitzer}, that reveal a hot dust shell with unusual emission features. ",
        "conclusions": "We have described the dust shell of the classical nova DZ~Cru, which shows several broad emission features. A number of these features remain unidentified. However the ``8.2'' feature in DZ~Cru and other novae may be associated with a UIR carrier in which aliphatic bonds predominate, likely as a consequence of the continued presence of dense clumps in the ejecta which shield the carrier from the hard radiation field of the nova. We anticipate that, as the ejecta (and clumps) disperse, a shift from aliphatic to aromatic may occur, accompanied by the shift of the UIR features to the blue; however it is more likely that the carrier will not survive exposure to the radiation. There are also implications for features in the 3\\mic-band, and it is unfortunate that spectra in this crucial region do not appear to have been obtained in the case of DZ~Cru.  Continued IR spectroscopy of DZ~Cru in the 3\\mic\\ and 7--13\\mic\\ regions would be valuable to observe the effects of dispersal on the UIR (and other) features."
    },
    "1005/1005.0411_arXiv.txt": {
        "abstract": "I provide notes on the NFW, Einasto, S\\'ersic, and other mass profiles which provide good fits to simulated dark matter halos (\\S\\ref{massprofiles}). I summarize various published $c(M)$ relations: halo concentration as a function of mass (\\S\\ref{S:cM}). The definition of the virial radius is discussed and relations are given to convert $c_{vir}$, $M_{vir}$, and $r_{vir}$ between various defined values of the halo overdensity (\\S\\ref{virial}). ",
        "introduction": "\\label{S:cM}  The mass profiles of galaxy clusters appear to be more centrally concentrated than realized in simulations \\citep{Broadhurst08,Oguri09,Sereno10}. If true, this may be evidence for Early Dark Energy \\citep[see e.g.,][]{GrossiSpringel09}. Or perhaps there is a less exciting explanation \\citep[e.g.,][]{BarkanaLoeb09,LapiCavaliere09,Meneghetti10}. For more details, see my discussion in \\cite{Coe10}. More conclusive results are expected from the CLASH HST MCT project\\footnote{Cluster Lensing And Supernova survey with Hubble, \\href{http://www.stsci.edu/~postman/CLASH/}{http://www.stsci.edu/$\\sim$postman/CLASH/} } and perhaps LoCuSS \\citep{Okabe09,Richard10}.  More massive halos generally have lower concentrations than less massive halos. This is seen in both simulations and observations, though less clearly so far in the latter (see below). More massive halos form later, resulting in lower concentrations reflecting the lower background density at the time of formation \\cite{NFW96}.  For a given radial mass profile (see \\S\\ref{massprofiles}), the concentration is defined as: \\\\ \\begin{equation} c _{vir} = r_{vir} / r_{-2}, \\end{equation} \\\\ a mishmash ratio of the virial radius $r_{vir}$ and the radius $r_{-2}$ at which $\\rho \\propto r^{-2}$. For an NFW profile, $r_{-2} = r_s$. The definition of the virial radius $r_{vir}$ is discussed at length in \\S\\ref{virial}, but it is typically approximated as the region within which there is an average overdensity of a certain value ($\\Delta_c \\sim100$ or 200) above $\\rho_{crit}$.  % For clarity, one may quote the exact value of $\\Delta_c$ used: $c_{200}$, for example.  In principle concentrations could be derived using any radial fitting profile (\\S\\ref{massprofiles}). However the choice does matter as the profiles behave differently between $r_{-2}$ and $r_{vir}$. Concentrations derived from NFW and Einasto fits to the same halos \\citep{Duffy08} are compared in Fig.~\\ref{cMzEinrel}. Einasto $c(M)$ relations have been derived for the Millennium simulation relaxed \\citep{Gao08} and all \\cite{HayashiWhite08} halos. Below we focus on $c(M)$ relations derived from NFW fits.   \\subsection{Current $c(M)$ measurements from NFW profile fits}  The current best estimates for $c(M,z)$ are probably those given by \\cite{Duffy08} and \\cite{Maccio08}. Their findings are similar. Both analyze simulations which use the WMAP5 cosmology, resulting in $\\sim 20\\%$ lower concentrations than WMAP1 (Table \\ref{tab:cosmo}) as used in the Millennium simulation \\citep{Neto07}, for example. \\cite{Duffy08} find that present-day ($z=0$) halos follow the following mass-concentration relation: \\\\ \\begin{equation} \\label{Duffy08} c_{200} \\simeq 5.74 \\left( \\frac{M_{200}}{2 \\times 10^{12} h^{-1} M_\\odot} \\right) ^ {-0.097}. \\end{equation} \\\\ They provide a separate relation for relaxed clusters which are more symmetric and thus better fit by radial profiles such as NFW. These have $15-20\\%$ higher concentrations (Fig.~\\ref{crelrat}): \\\\ \\begin{equation} \\label{Duffy08} c_{200} \\simeq 6.67 \\left( \\frac{M_{200}}{2 \\times 10^{12} h^{-1} M_\\odot} \\right) ^ {-0.092}. \\end{equation} \\\\ Intrinsic scatters are $\\Delta \\log_{10} (c_{200}) \\simeq 0.15$. These relations are plotted in Fig.~\\ref{cMDuffyMaccio} along with corresponding relations from \\cite{Maccio08}.  \\cite{Duffy08} also supply fitted functions to halos spanning the redshift range $z = 0 - 2$. Full: \\begin{equation} \\label{Duffy08} c_{200} \\simeq \\frac{5.71}{(1+z)^{0.47}} \\left( \\frac{M_{200}}{2 \\times 10^{12} h^{-1} M_\\odot} \\right) ^ {-0.097} \\end{equation} \\\\ and relaxed: \\\\ \\begin{equation} \\label{Duffy08} c_{200} \\simeq \\frac{6.71}{(1+z)^{0.44}} \\left( \\frac{M_{200}}{2 \\times 10^{12} h^{-1} M_\\odot} \\right) ^ {-0.092}. \\end{equation} \\\\ In their Table 1, they provide uncertainties for these fit parameters as well as corresponding values for $c_{vir}$ and $M_{vir}$. These $c_{200}(M_{200},z)$ relations are plotted in Figs.~\\ref{cMzM} \\& \\ref{cMzz}. Also plotted are the corresponding $c_{vir}(M_{vir},z)$ relations provided by \\citep{Duffy08}. In Fig.~\\ref{cMzz}, we plot the \\cite{Bullock01} $c \\propto (1+z)^{-1}$ scaling for comparison. Note that Duffy08 find weaker dependencies on redshift: $c_{200} \\propto (1+z)^{-0.45}$; $c_{vir} \\propto (1+z)^{-0.70}$.  Halo concentrations are sensitive to cosmology. A higher $\\sigma_8$ causes halos to form earlier, resulting in higher concentrations. This was the case in the Millennium simulations which used the WMAP 1-year cosmology, including $\\sigma_8 = 0.9$. This yields concentrations $\\sim 20\\%$ higher than found in simulations which use WMAP5's $\\sigma_8 = 0.796$\\footnote{This value is in excellent agreement with the WMAP 7-year maximum likelihood value $\\sigma_8 = 0.803$ \\citep{WMAP7}.} \\citep{Duffy08}. The effect of cosmology was explored in more detail by \\cite{Maccio08}. These effects are shown in Fig.~\\ref{cMcosmo}.  Various derived $c(M)$ relations (for $z=0$) are plotted in Fig.~\\ref{cMcompare}. In our Tables \\ref{tab:cM}, \\ref{tab:cMvir}, and \\ref{tab:cMz}, we provide $c_{200}(M_{200})$, $c_{vir}(M_{vir})$, and $c(M,z)$ relations, respectively, as derived by \\cite{Duffy08}, \\cite{Maccio08}, \\cite{Neto07}, \\cite{Bullock01}, \\cite{Hennawi07}, and \\cite{Gentile07}.\\footnote{I'll take a crack at completeness, and mention other papers with $c(M)$ relations neglected here (for no particular reason): \\cite{Eke01,Wechsler02,Alam02,Zhao03b,Dolag04,Kuhlen05,Lu06,Shaw06,Dutton07,Gnedin07}. And I have probably missed still others!} The latter is the original NFW $c(M)$ prescription updated to the WMAP3 cosmology.  The various simulations considered here are outlined in Table~\\ref{tab:sims}. The relevant details of their adopted cosmologies ($\\Omega_m$, $\\sigma_8$) are given in Table~\\ref{tab:cosmo}. We provide the range of halo masses produced in each simulation. The dangers of extrapolating $c(M)$ relations beyond these ranges have been cited by \\cite{Zhao03b}, for example.  \\citet[their Fig.~5]{Neto07} find that 10,000 particles within the virial radius are required to yield robust concentration measurements. They note that using fewer particles introduces scatter but does not appear to introduce bias in their concentration measurements.   \\cite{Hennawi07} measure significantly larger concentrations for galaxy clusters in their simulations. Their cluster concentrations are $\\sim 50\\%$ and $\\sim 80\\%$ larger than found by \\cite{Duffy08} and \\cite{Maccio08}, respectively (Fig.~\\ref{cMcompare}, {\\it right}). Their use of $\\sigma_8 = 0.95$ probably only results in concentrations inflated by $\\sim 20\\%$ compared to the WMAP5 $\\sigma_8 = 0.796$ simulations. The remaining disagreement may be a result of their halo density fitting procedure which they claim is better for comparison with lensing measurements. Specifically, they assign large uncertainties to radial bins with large subhalos. This may bias the fitted profiles to be low at large radius (where large subhalos typically reside) resulting in higher concentrations. These results may considerably ease tensions between observed and simulated halo concentrations. The differences in fitting procedure should be better studied and understood.   \\subsection{Care in citing concentration expectations}  A concern often noted is that the concentration measured for A1689 (in every study to date) is higher than that found in simulations for a halo of A1689's mass. The concentration found in simulations has been cited loosely as $c \\sim 5$ or $c \\sim 5.5$ using a relation given by \\cite{Bullock01}: \\\\ \\begin{equation} \\label{NFWcBullock01} c_{vir} \\simeq \\frac{9}{1+z} \\left( \\frac{M_{vir}}{1.3 \\times 10^{13} h^{-1} M_\\odot} \\right) ^ {-0.13}, \\end{equation} But the expected concentration is actually lower ($c_{200} \\sim 3.0$, exacerbating the disagreement with observations) for four reasons: \\vspace{-0.08in} \\begin{itemize} \\item $c_{200} < c_{vir}$ \\vspace{-0.08in} \\item $z > 0$ \\vspace{-0.08in} \\item WMAP7 vs.~WMAP1 \\vspace{-0.08in} \\item $M_{vir} \\approx 1.4 \\times 10^{15} M_\\odot / h > 10^{15} M_\\odot$ \\end{itemize} \\vspace{-0.05in} We also note that \\cite{Bullock01} did not simulate halos as massive as A1689, with their most massive halos weighing in at $M \\sim 10^{14} M_\\odot / h$.   \\subsection{The observed $c(M)$ relation}  Based on a compilation of 62 published measurements (including 10 new measurements) of halo $c_{vir}$ and $M_{vir}$, \\cite{ComerfordNatarajan07} find the following relation (with a large scatter): \\\\ \\begin{equation} \\label{NFWcComerford07} c_{vir} \\simeq \\frac{14.5 \\pm 6.4}{1+z} \\left( \\frac{M_{vir}}{1.3 \\times 10^{13} h^{-1} M_\\odot} \\right) ^ {-0.15 \\pm 0.13}. \\end{equation} \\\\ For clusters as massive as A1689, Comerford's relation converges toward that of \\cite{Hennawi07}, the former being only slightly higher.  This and other observed $c(M)$ relations are shown in Fig.~\\ref{cMcompare} and detailed in Table~\\ref{tab:cMobs}. The \\cite{ComerfordNatarajan07} compilation includes both lensing and X-ray determinations of $c$ and $M$, including the X-ray samples presented by \\cite{Buote07} and \\cite{SchmidtAllen07}. Each of these papers presented their own $c(M)$ relation. A recent $c(M)$ relation from weak lensing of individual halos was presented by \\cite{Okabe09}. And $c(M)$ derived from {\\it stacked} weak lensing analyses were presented by \\cite{Johnston07} and \\cite{Mandelbaum08}. It seems apparent that one should study a wide enough range of halo masses to obtain a confident $c(M)$ relation.   \\vspace{0.2in} ",
        "conclusions": ""
    },
    "1005/1005.4250.txt": {
        "abstract": "Recent measurements of hot and cold spots on the cosmic microwave background (CMB) sky suggest a presence of super-structures on ($>100\\,h^{-1}$Mpc) scales.  We develop a new formalism to estimate the expected amplitude of temperature fluctuations due to the integrated Sachs-Wolfe (ISW) effect from prominent quasi-linear structures. Applying the developed tools to the observed ISW signals from voids and clusters in catalogs of galaxies at redshifts $z<1$, we find that they indeed imply a presence of quasi-linear super-structures with a comoving radius $100\\sim 300\\,h^{-1}$Mpc and a density contrast $|\\delta|\\sim O(0.1)$. We also find that the observed ISW signals are at odd with the concordant $\\Lambda$ cold dark matter (CDM) model that predicts Gaussian primordial perturbations at $\\gtrsim 3\\, \\sigma$ level. We confirm that the mean temperature around the CMB cold spot in the southern Galactic hemisphere filtered by a compensating top-hat filter deviates from a mean value at $\\sim 3\\,\\sigma$ level, implying that a quasi-linear supervoid or an underdensity region surrounded by a massive wall may reside at low redshifts $z<0.3$ and the actual angular size ($16^{\\circ}-17^\\circ$) may be larger than the apparent size ($4^\\circ-10^\\circ$) discussed in literature. Possible solutions are briefly discussed. ",
        "introduction": "Although the $\\Lambda$ cold dark matter (CDM) models have succeeded in explaining a number of observations, some problems remain unresolved. For example, origins of a possible break of statistical isotropy in the large-angle cosmic microwave background (CMB) anisotropy (Tegmark et al. 2003; Eriksen et al. 2004; Vielva et al. 2004) and a possible discrepancy between observed and theoretically predicted galaxy-CMB cross-correlation (Rassat et al. 2007; Ho et al. 2008) are still not understood well. These observational results imply that structures on scales larger than $\\gtrsim 100$ Mpc (super-structures) in our local universe are more lumpy than expected (Afshordi et al. 2009).  As the origin of the large-angle CMB anomalies, many authors have considered a possibility that the CMB is affected by local inhomogeneities (Moffat 2005; Tomita 2005a,b; Cooray \\& Seto 2005; Rakic \\& Schwartz 2007). Inoue \\& Silk (2006, 2007) have shown that a particular configuration of compensated quasi-linear supervoids can explain most of the features of the anomalies. Subsequent theoretical analyses have shown that the CMB temperature distribution for quasi-linear structures can be skewed toward low temperature due to the second order integrated Sachs-Wolfe (ISW)(or Rees-Sciama) effect (Tomita \\& Inoue 2008; Sakai \\& Inoue 2008).  In fact, Granett et al. (2008) found a significant ISW signal at the scale of $4^\\circ-6^\\circ$ at redshifts around $z \\sim 0.5$ and a weak signal of negatively skewed temperature distribution for distinct voids and clusters at redshifts $0.4<z<0.75$. Moreover, Francis \\& Peacock (2009) have shown that the ISW effect due to local structures at redshift $z<0.3$ significantly affects the large-angle CMB anisotropies and that some of the CMB anomalies no longer persist after subtraction of the ISW contribution.  These observations of galaxy-CMB cross-correlation may suggest an existence of anomalously large perturbations or new physics on scales $>100$ Mpc.  In order to evaluate the significance of the ISW signals for prominent structures, N-body simulations on cosmological scales seem to be suitable(Cai et al. 2010) for this purpose. However, the computation time is relatively long and finding physical interpretation from a number of numerical results is sometimes difficult. In contrast, analytical methods are suitable for estimating the order of statistical significance in a relatively short time, and physical interpretations are often simpler.  In this paper, we evaluate the statistical significance of the ISW signals for prominent super-structures based on an analytic method and try to construct simple models that are consistent with the data. In section 2, we develop a formalism for analytically calculating the ISW signal due to prominent non-linear super-structures based on a spherically symmetric homogeneous collapse model and we study the effect of non-linearity and inhomogeneity of such structures.  In section 3, we apply the developed method to observed data and calculate the statistical significance of the discrepancy between the predicted and the observed ISW signals. In section 4, we discuss the origin of the observed discrepancy. In section 5, we summarize our results and discuss some unresolved issues. In the following, unless noted, we assume a concordant $\\Lambda$CDM cosmology with $(\\Omega_{m,0}, \\Omega_{\\Lambda,0}, \\Omega_{b,0}, h,\\sigma_8,n)=(0.26,0.74,0.044,0.72,0.80,0.96)$, which agrees with the recent CMB and large-scale structure data (S\\'anchez et al. 2009). \\pagebreak ",
        "conclusions": "In this paper, we have shown that recent observations imply a presence of quasi-linear super-structures with a comoving radius $100-300\\,h^{-1}$Mpc at redshifts $z<1$.  Observations are at odd with the concordant $\\Lambda$CDM cosmology that predicts Gaussian primordial perturbations at $>3\\,\\sigma$ level.  First, we have developed a formalism to estimate the amplitude of the ISW signal for prominent structures based on thin-shell approximation and the homogeneous collapse model. From comparison with other calculations based on perturbation theory and the LTB solution, we have found that our simple model works well for estimating the ISW signal for quasi-linear superstructures in $\\Lambda$-dominated universes.   Secondly, we have applied our developed tools to observations of the ISW signals using the SDSS-LRG catalog, the 2MASS catalog, and the cold spot in the Galactic southern hemisphere in the WMAP data. The ISW signals from stacked images for the SDSS-LRG catalog is inconsistent with the predicted values in the concordant $\\Lambda$CDM model at more than $3\\, \\sigma$. The radial profiles of the stacked image show a hot-ring around a cold spot for voids and a dip at the center of a hot-spot without a cold-ring for clusters. These non-linear features are also reproduced by our models using the LTB solutions although the agreement is not perfect. The asymmetrical features suggest that the observed super-structures are in quasi-linear regime rather than linear regime. The amplitudes of the ISW signals obtained from the 2MASS catalog at redshifts $0.1<z<0.3$ are found to be several times larger than expected values. We have confirmed that the mean temperature around the cold spot filtered by a compensating top-hat filter with angular scale $\\theta_{out}=16^\\circ-17^\\circ$ deviates from the mean value at roughly $3\\,\\sigma$ level suggesting a presence of a hot-ring around the cold spot. Note that our finding is consistent with the previous result that the cold spot itself is not unusual but the hot-ring plus the inner cold region is found to be unusual (Zhang \\& Huterer 2009). This implies that a supervoid may reside at low redshift $z<0.3$ and the angular size may be larger ($=16^\\circ-17^\\circ$) than considered in literature (Masina \\& Notari 2009).  Finally, we have discussed possible causes of the discrepancy between the theory and observation, namely, observational systematics, primordial non-Gaussianity, features in power spectrum, dark energy or alternative gravity.  We have not considered effects of non-spherical collapse which are important for improving estimation of the mass function of non-linear objects and effects of uncompensated mass distribution for super-structures. The extension of our analysis to more elaborate ones incorporating these effects would be helpful for realizing detailed comparison between the theory and the observation.   Future surveys of the CMB, galaxy distribution, weak lensing and theoretical studies on dark energy/alternative gravity and inhomogeneous cosmology will certainly yield fruitful results for solving the puzzle. \\\\ \\\\"
    },
    "1005/1005.0627_arXiv.txt": {
        "abstract": "Black holes of mass $M$ must have a spin angular momentum $S$ below the Kerr limit $(\\chi \\equiv S/M^2 \\leq 1)$, but whether astrophysical black holes can attain this limiting spin depends on their accretion history.  Gas accretion from a thin disk limits the black-hole spin to $\\chi_{\\rm gas} \\lesssim 0.9980 \\pm 0.0002$, as electromagnetic radiation from this disk with retrograde angular momentum is preferentially absorbed by the black hole.  Extrapolation of numerical-relativity simulations of equal-mass binary black-hole mergers to maximum initial spins suggests these mergers yield a maximum spin $\\chi_{eq} \\lesssim 0.95$.  Here we show that for smaller mass ratios $q \\equiv m/M \\ll 1$, the superradiant extraction of angular momentum from the larger black hole imposes a fundamental limit $\\chi_{\\rm lim} \\lesssim 0.9979 \\pm 0.0001$ on the final black-hole spin even in the test-particle limit $(q \\to 0)$ of binary black-hole mergers.  The nearly equal values of $\\chi_{\\rm gas}$ and $\\chi_{\\rm lim}$ imply that measurement of supermassive black-hole spins cannot distinguish a black hole built by gas accretion from one assembled by the gravitational inspiral of a disk of compact stellar remnants.  We also show how superradiant scattering alters the mass and spin predicted by models derived from extrapolating test-particle mergers to finite mass ratios. ",
        "introduction": "\\label{S:intro}  Supermassive black holes (SBHs) reside in the centers of most large galaxies.  While a few nearby SBHs can be detected by their gravitational influence on surrounding stars, the majority of SBHs are observed electromagnetically as active galactic nuclei (AGN).  The same accretion flows that release energy to power AGN also supply energy and angular momentum to the black holes themselves increasing their mass $M$ and spin $S$ \\footnote{Throughout this paper we express physical quantities in units where Newton's constant $G$ and the speed of light $c$ are unity, in which case the spin angular momentum $S$ has dimensions of $M^2$.}.  By definition, black holes possess an event horizon from within which nothing can escape to future null infinity.  Black holes described by the Kerr metric \\cite{Kerr:1963ud} only possess an event horizon for $\\chi \\equiv S/M^2 \\leq 1$, setting a fundamental upper limit on a black hole's possible spin.  Whether or not astrophysical SBHs saturate or even exceed this Kerr limit is an important test of general relativity.  Black-hole spins also probe their assembly history.  SBHs grow both by gas accretion and mergers driven by the gravitational inspiral of binary companions.  These two growth mechanisms may supply mass and angular momentum to the SBHs in different ratios, allowing measurements of black-hole spin to distinguish between them. Reverberation mapping of the iron K$\\alpha$ line in AGN x-ray spectra has been proposed as just such a means of measuring SBH spins \\cite{Reynolds:1998ie}.  XMM-Newton observations analyzed with this technique have been used to constrain the spin of the SBH hosted by the Seyfert 1.2 galaxy MCG-06-30-15 to $\\chi = 0.989_{-0.002}^{+0.009}$ at 90\\% confidence \\cite{Brenneman:2006hw}. A closer examination of how SBHs acquire their spins is thus of both theoretical and observational importance.  According to Bardeen \\cite{Bardeen:1970}, a nonspinning black hole can attain the maximum Kerr spin $\\chi = 1$ after accreting a finite mass of test particles freely falling from its innermost stable circular orbit (ISCO).  However, material on circular orbits with radii greater than that of the ISCO cannot be accreted unless it has some mechanism to shed its excess angular momentum.  In a standard geometrically thin, optically thick accretion disk \\cite{Shakura:1972te}, this mechanism is viscous stress within the disk that also heats it locally and transports energy outwards.  This heating produces an energy flux at the disk's surface that will be radiated in all directions.  A small fraction of these radiated photons do not escape to infinity but are instead captured by the black hole itself.  Photons with negative angular momentum with respect to the black hole's spin have a larger capture cross-section than those with positive angular momentum \\cite{Godfrey:1970am}, implying that the accreted radiation will counteract the directly advected material which always acts to spin the hole up for $\\chi < 1$.  These two sources of angular momentum cancel for black-hole spins $\\chi_{\\rm gas} \\simeq 0.998$, with black holes of greater spins spinning down to this value after accreting a mass $m \\simeq 0.05M$ \\cite{Thorne:1974ve}.  Different accretion flows will supply energy and angular momentum to the black hole in different ratios, altering the value of the limiting spin $\\chi_{\\rm gas}$.  Advection-dominated accretion flows (ADAFs) do not cool efficiently, and therefore a fraction $f > 0$ of the gravitational energy dissipated prior to the ISCO is advected by the black hole rather than radiated to infinity.  Viscous stresses in these flows proportional to the Shakura-Sunyaev parameter $\\alpha$ \\cite{Shakura:1972te} also reduce the specific angular momentum of the accreted material below its value in the vacuum at the ISCO.  The former effect increases $M$ in the denominator in the definition $\\chi \\equiv S/M^2$ of the dimensionless spin, while the latter effect reduces $S$ in the numerator \\cite{Popham:1998iw}.  The magnetorotational instability \\cite{Balbus:1991ay} also torques gas at the ISCO, and in addition can launch jets which further limit the spin as shown in magnetohydrodynamic simulations \\cite{Gammie:2003qi}.  Analytic fits to the ADAF simulations of Popham and Gammie \\cite{Popham:1998iw} suggest that black holes can spin up to $\\chi_{\\rm ADAF} \\simeq 0.96$, and that the inclusion of jets calibrated by the magnetohydrodynamic simulations of Gammie {\\it et al.} \\cite{Gammie:2003qi} reduces this limit to $\\chi_{\\rm jets} \\simeq 0.93$ \\cite{Benson:2009kx}.  We see that limits on black-hole spin depend greatly on the nature of the accretion flow, and that no model-independent constraints exclude the limit $\\chi_{\\rm gas} \\simeq 0.998$ for thin disks first obtained by Thorne \\cite{Thorne:1974ve}.  The maximum spin of black holes produced in binary mergers is also greatly uncertain.  The most accurate numerical-relativity simulations have been of equal-mass $q \\equiv m_2/m_1 = 1$ black holes with spins $\\chi_1 = \\chi_2$ aligned or antialigned with the orbital angular momentum.  These simulations found $\\chi = 0.68646 \\pm 0.00004$ for nonspinning binaries \\cite{Scheel:2008rj}, and $\\chi = 0.547812 \\pm 0.000009$ for equal-mass binaries with initial spins $\\chi_1 = \\chi_2 = -0.43757$ \\cite{Chu:2009md}.  The spins of such aligned binaries do not precess prior to merger, removing one complication in determining their magnitude and direction.  We will therefore restrict our analysis to aligned binaries in this paper, though with modest additional work it could be extended to binaries of arbitrarily oriented spins.  Although these numerically determined spins are remarkably precise, two significant obstacles still prevent the determination of final black-hole spins from more generic mergers.  The first is the inability of the commonly used conformally flat Bowen-York initial data to adequately approximate initial binary black holes with spins greater than $\\chi_i \\simeq 0.93$ \\cite{Lovelace:2008tw}.  An extrapolation of simulations with $\\chi_i \\leq 0.9$ suggests that the merger of equal-mass black holes with maximal initial spins aligned with their orbital angular momentum yields a final black hole with spin $\\chi = 0.951 \\pm 0.004$ \\cite{Marronetti:2007wz}.  However, it remains to be proven that such a smooth extrapolation in initial spin holds all the way to the Kerr limit.  The second obstacle to determining spins from generic mergers is the increased computational resources needed to simulate mergers with mass ratios $q \\leq 1$.  This increased computational demand results from the need for more closely spaced numerical grid points to resolve the horizon of the smaller black hole, shorter timesteps owing to the shorter light propagation time between grid points, and longer simulations to capture the larger number of orbits per unit increase in orbital frequency prior to merger.  The smallest mass ratio that has currently been simulated is $q = 0.01$ \\cite{Lousto:2010ut}; only a handful of simulated mergers with $q \\leq 0.1$, all of initially nonspinning black holes, have been published \\cite{Gonzalez:2008bi,Lousto:2010tb,Lousto:2010qx}.  New numerical techniques will be required to make much progress beyond this point.  Relativists have attempted to surmount these obstacles by inventing fitting formulas that are functions of $q$, $\\chi_1$, and $\\chi_2$, calibrating the coefficients in these formulas with existing simulations, and then extrapolating them to higher spins and lower mass ratios than can currently be simulated.  This approach is very effective in the region of the parameter space $\\{q, \\chi_1, \\chi_2 \\}$ near the simulations with which the fitting formulas were calibrated, but can break down outside this region.  The probable reason for this is that all three parameters listed above equal unity for equal-mass, maximally spinning, aligned mergers, and thus polynomials in these parameters will converge slowly if at all.  Even the symmetric mass ratio $\\eta \\equiv m_1m_2/(m_1 + m_2)^2 = q/(1 + q)^2$ gets as large as 1/4 for equal-mass mergers, suggesting that polynomials in this parameter will not converge quickly either.  At the time this paper was written, simulations were restricted to the region $0.1 \\leq q \\leq 1$ ($0.0826 \\leq \\eta \\leq 0.25$) which provided a short lever arm over which to calibrate terms with different $q$ dependence.  The degeneracy between terms makes predicting final spins in the test-particle ($q \\to 0$) limit very uncertain.  For example, several fitting formulas predict final spins above the Kerr limit for maximally spinning aligned mergers with mass ratios as modest as $q \\lesssim 0.25$.  The publication of a simulation with $q = 0.01$ \\cite{Lousto:2010ut} during the preparation of this paper shows the rapid progress towards numerical simulations of extreme-mass-ratio mergers, but much work remains before such simulations are available for generically spinning black holes.  Fortunately our analytical knowledge of the geodesics of the Kerr metric can help us understand binary black-hole mergers in the test-particle limit.  These mergers can be broken down into three stages: inspiral, plunge, and ringdown.  During the inspiral stage, the orbit of the test particle adiabatically evolves through a series of geodesics with successively lower energy and angular momentum as these quantities are radiated away through the emission of gravitational waves.  This gravitational radiation circularizes initially eccentric orbits while the evolution is post-Newtonian \\cite{Peters:1963ux}, implying that for many though possibly not all astrophysical mergers the orbit will have circularized before the test particle reaches the ISCO.  After this point, during the brief plunge stage, the test particle rapidly falls into the event horizon of the larger black hole.  Finally, the quasinormal modes excited during the merger ring down as the newly formed black hole settles into its final Kerr configuration.  In the test-particle limit, the energy and angular momentum radiated during the inspiral stage scales linearly with $q$, while that radiated during plunge and ringdown scales as a higher power in the small parameter $q$.  This suggests that one can reasonably predict the mass and spin of the final black hole by equating these quantities to the energy and total angular momentum of the binary at the ISCO.  This prediction assumes that during the inspiral all of the gravitational waves are radiated outwards to infinity rather than downwards to the larger black hole's event horizon.  For highly spinning black holes, these downward gravitational waves will be superradiantly scattered, extracting energy and angular momentum from the larger black hole.  Individual modes can be amplified as much as 138\\% by a maximally spinning black hole \\cite{Teukolsky:1974yv}; the total energy flux radiated to infinity by a test particle at the ISCO will be amplified by 12.9\\% for a black hole of spin $\\chi = 0.999$ \\cite{Finn:2000sy}.  Although this amplification decreases rapidly as a function of the orbital radius, the total energy and angular momentum extracted throughout the inspiral still scale linearly in $q$ and will therefore remain the dominant factor affecting the final mass and spin after the energy and angular momentum advected with the test particle itself during merger.  The primary goal of this paper is to determine how much the superradiant scattering of gravitational waves during the inspiral reduces the maximum spin achievable by binary black-hole mergers below the Kerr limit first predicted by Bardeen \\cite{Bardeen:1970}.  In Sec.~\\ref{S:test} we will review the previous analytic predictions for the final mass and spin that serve as our starting point.  Our method for calculating the energy and angular momentum extracted from the spinning black hole will be described in Sec.~\\ref{S:super}.  The results of this calculation are presented in Sec.~\\ref{S:res}.  In Sec.~\\ref{S:num} we compare our model to others in the literature, extrapolate it to comparable-mass mergers and compare with numerical-relativity simulations, and suggest how it might be combined with fitting formulas to produce an approximation valid over the entire domain $0 \\leq q \\leq 1$.  A brief summary and a few final remarks on the astrophysical implications of our analysis are offered in Sec.~\\ref{S:disc}. ",
        "conclusions": "\\label{S:disc}  The superradiant scattering of gravitational waves emitted during the inspiral sets a fundamental upper limit $\\chi_{\\rm lim} = 0.9979 \\pm 0.0001$ on the spin a black hole may attain by accreting test particles on quasicircular equatorial orbits.  For this limiting spin, the energy and angular momentum advected when the test particle falls from the ISCO combines with that extracted by superradiant scattering to leave the spin $\\chi = S/M^2$ of the black hole unchanged.  Black-hole spins greater than $\\chi_{\\rm lim}$ will be reduced by mergers with test particles on quasicircular equatorial orbits, even though the physical angular momentum $S$ and mass $M$ will both increase.  Jacobson and Sotiriou \\cite{Jacobson:2009kt} recently argued that small but finite-mass particles on highly {\\it hyperbolic} orbits could spin black holes above $\\chi_{\\lim}$ and even above the Kerr limit itself.  However, this study neglected gravitational radiation and any superradiant enhancement, which could be considerable for these highly spinning mergers.  Such hyperbolic orbits are also extremely unlikely in any realistic astrophysical scenario.  Our limit is nearly indistinguishable from the limit $\\chi_{\\rm gas} = 0.9980 \\pm 0.0002$ for black holes grown through thin-disk gas accretion \\cite{Thorne:1974ve}, suggesting that spin measurements alone cannot determine the relative contribution of these two channels of black-hole growth.  If gas accretion occurs through an advection-dominated accretion flow (ADAF) rather than a thin disk, the limiting black-hole spin is reduced to $\\chi_{\\rm ADAF} \\simeq 0.96$, and the launching of jets can further reduce this spin to $\\chi_{\\rm jets} \\simeq 0.93$ \\cite{Benson:2009kx}.  Since our limit $\\chi_{\\rm lim}$ is larger than for other modes of black-hole growth, it seems to be a fairly robust upper bound on spins that can be obtained by {\\it any} astrophysical means.  Although SBHs are grown primarily through gas accretion rather than test-particle mergers, our limiting spin $\\chi_{\\rm lim}$ can still be astrophysically relevant.  Gas accretion occurs episodically, when galactic major mergers produce global torques that funnel large amounts of gas to galactic centers \\cite{Mihos:1995ri,Barnes:1996qt}. Accretion of this gas onto the SBH leads to AGN feedback \\cite{Silk:1997xw} which simulations demonstrate is capable of suppressing further accretion and black-hole growth \\cite{DiMatteo:2005sd}.  By contrast, compact objects such as white dwarfs, neutron stars, and stellar-mass black holes (effectively test particles when orbiting an SBH) will inspiral continuously as they are scattered into the ``loss cone'' of orbits for which gravitational inspiral occurs more rapidly than subsequent scattering \\cite{Frank:1976uy}.  Although the same global torques driven by galactic major mergers may enhance the rate of these extreme-mass-ratio inspirals (EMRIs), other dynamical processes will refill the loss cone even between periods of AGN activity \\cite{Sigurdsson:2003ei}.  Event rates for EMRIs depend sensitively on highly uncertain factors such as the dynamical state, star-formation history, binary fraction, and initial mass function within nuclear star clusters.  The capture rate for $10 M_\\odot$ black holes by a $10^6 M_\\odot$ SBH can be as high as $10^{-4} {\\rm yr}^{-1}$ immediately following a nuclear starburst, and could average $10^{-6} {\\rm yr}^{-1}$ for nucleated spiral galaxies like the Milky Way \\cite{Sigurdsson:1996uz}.  If such a galaxy has not had a major merger since redshift $z \\simeq 1$ (about 7.5 Gyr ago), its SBH would aqcuire a fraction \\begin{equation} \\label{E:EMRIf} \\frac{\\Delta M}{M_0} \\simeq \\frac{10 M_\\odot \\times 10^{-6}{\\rm yr}^{-1} \\times 7.5~{\\rm Gyr}}{10^6 M_\\odot} = 0.075 \\end{equation} of its mass through EMRIs.  We see from Fig.~\\ref{F:evochi} that this is a large enough fraction to drive the SBH spin towards our limiting value $\\chi_{\\rm lim}$ from the initial spins $\\chi_0 \\gtrsim 0.9$ expected from gas accretion.  This suggests that EMRIs can potentially affect SBH spin evolution, though modest gas accretion can also occur between major mergers and not all EMRIs will be on circular, equatorial orbits.  Even if our limiting spin $\\chi_{\\rm lim}$ can be realized astrophysically, does its marginal difference from the Kerr limit have any observable consequences?  The binding energy per unit mass at the ISCO $1 - E_{\\rm ISCO}$ depends very sensitively on $\\chi$ near the Kerr limit, and sets an upper bound on the AGN efficiency $\\varepsilon \\equiv L/\\dot{M}c^2$, where $L$ is the AGN luminosity, $\\dot{M}$ is the accretion rate, and $c$ is the speed of light.  Soltan \\cite{Soltan:1982vf} argued that the total cosmological mass in SBHs could be estimated by equating the observed energy emitted by AGN to that released during theoretical models of SBH growth.  If spins at the Kerr limit were used for an analysis of this kind instead of our limiting spin $\\chi_{\\rm lim} = 0.9979$, the AGN efficiency $\\varepsilon$ would be overestimated by $E_{\\rm ISCO}(\\chi_{\\rm lim}) - E_{\\rm ISCO}(1) = 0.1031$ leading to a corresponding underestimate in the total SBH mass.  In addition to their influence on the SBH spin, the EMRIs themselves are an important source for gravitational-wave detectors such as LIGO and LISA.  Decreasing the spin from the Kerr limit to $\\chi_{\\rm lim}$ moves the ISCO radius out from $r_{\\rm ISCO}(1) = m_1$ to $r_{\\rm ISCO}(\\chi_{\\rm lim}) = 1.242 m_1$, and decreases the ISCO orbital frequency from $\\Omega_{\\rm ISCO}(1) = 0.50 m_{1}^{-1}$ to $\\Omega_{\\rm ISCO}(\\chi_{\\rm lim}) = 0.42 m_{1}^{-1}$.  This decrease could be observable in high signal-to-noise EMRIs seen by both these experiments.  Our limiting spin $\\chi_{\\rm lim}$ is thus not only an interesting consequence of general relativity, but also one that has potentially observable implications for astrophysics and gravitational-wave detection.  \\vspace{0.3cm}  {\\bf Acknowledgements.}  We thank Scott Hughes for generously allowing us to use a version of his GREMLIN (Gravitational Radiation in the Extreme Mass ratio LIMit) code to calculate the energy and angular momentum radiated during the inspiral.  We also thank Andrew Benson, Chris Hirata, Pedro Marronetti, Uli Sperhake, and Kip Thorne for useful conversations.  The authors acknowledge support from NASA BEFS Grant No.~NNX07AH06G (PI:Phinney).  M. Kesden also acknowledges support from NSF Grant No.~PHY-0601459 (PI: Thorne)."
    },
    "1005/1005.2287_arXiv.txt": {
        "abstract": "During post-Newtonian evolution of a compact binary, a mass ratio~$\\nu $ different from 1 provides a second small parameter, which can lead to unexpected results. We present a \\textit{statistics of supermassive black hole candidates}, which enables us first to derive their mass distribution, then to establish a logarithmically even probability in $\\nu $ of the mass ratios at their encounter. In the mass ratio range $\\nu \\in \\left( 1/30,1/3\\right) $ of supermassive black hole mergers representing $40\\%$ of all possible cases, the combined effect of spin-orbit precession and gravitational radiation leads to a spin-flip of the dominant spin during the \\textit{inspiral} phase of the merger. This provides a mechanism for explaining a large set of observations on X-shaped radio galaxies. In another $40\\%\\,$ with mass ratios $\\nu \\in \\left( 1/30,1/1000\\right) $ a spin-flip never occurs, while in the remaining $20\\%$ of mergers with mass ratios $\\nu \\in \\left( 1/3,1\\right) $ it may occur during the plunge. We analyze the magnitude of the spin-flip angle occurring during the inspiral as function of the mass ratio and original relative orientation of the spin and orbital angular momentum. We also derive a formula for the \\textit{final spin }at the end of the inspiral in this mass ratio range. ",
        "introduction": "During galaxy mergers, following a regime of slow approach due to dynamical friction, eventually the central supermassive black holes (SMBHs) approach each other to a separation of the order of $10^{3}$ Schwarzschild radii, when gravitational radiation takes over as the leading order dissipative effect. The Laser Interferometer Space Antenna (LISA, see \\cite{LISA}) is expected to detect merging binary SMBHs with masses $m_{1}+m_{2}\\le 10^{7}\\,$% solar masses (M$_{\\odot }$) up to redshift $z\\approx 30$. A post-Newtonian approach is well suited to describe their forthcoming \\textit{inspiral}, a regime we define in terms of the post-Newtonian (PN) parameter $\\varepsilon =Gm/c^{2}r$ $\\approx v^{2}/c^{2}\\in \\left( \\varepsilon _{in}=10^{-3},\\varepsilon _{fin}=10^{-1}\\right) $, where $r$ and $v$ characterize\\ the orbital separation (from the center of mass) and speed of the reduced mass particle, $G$ is the gravitational constant and $c$ the speed of light. Various corrections to the conservative dynamics add up to 2 PN, while the gravitational radiation results in dissipation of energy, angular momentum and orbital angular momentum at 2.5 PN.  The leading order conservative correction to the Newtonian dynamics in a compact binary, which results in a change of the orbital plane (defined by the direction $\\mathbf{\\hat{L}}_{N}$ of the Newtonian orbital angular momentum $\\mathbf{L}_{N}=\\mu \\mathbf{r}\\times \\mathbf{v}$ of the reduced mass particle $\\mu $) is the spin-orbit (SO) interaction \\cite{SO}, \\cite% {ACST}. The precessional time-scale (the time during which the normal to the orbit $\\mathbf{\\hat{L}}_{N}$ undergoes a full rotation) is longer than the orbital period, however shorter than the characteristic time-scale of gravitational radiation (defined as $L/\\dot{L}$, where $L$ is the magnitude of the total orbital angular momentum). Combined with the leading order gravitational radiation backreaction averaged over one quasicircular orbit, the SO correction provides a fair approximation to orbital dynamics, explored in Refs. \\cite{ACST}, \\cite{SpinFlip}.  X-shaped radio galaxies (XRGs) exhibit two pairs of radio lobes and jets \\cite{LeahyWilliams}, \\cite{Xshape}. A recent review \\cite{GBGW} summarizes the four different models explaining XRGs: galaxy harbouring twin AGNs, back-flow diversion models, rapid jet reorientation models, finally a new jet-shell interaction model. A large subset of the observations (excepting cases, when the jets are aligned with the optical axes of the host ellipticals \\cite{Battye}) are well-explained by the jet reorientation model, which in turn implies a spin-flip \\cite{LeahyWilliams}, \\cite% {Xspinflip} of the dominant black hole.  The details of how this would occur were worked out in Ref. \\cite{SpinFlip}. A key element was the determination of the typical mass ratio at SMBH mergers by a series of estimates, which resulted in mass ratios $\\nu =m_{2}/m_{1}=1/30\\,$\\ to $\\nu =1/3$. Because the spin scales with the mass squared, the second spin was neglected and only the dominant spin $\\mathbf{S}% _{\\mathbf{1}}$ (with magnitude $S_{1}$) kept. We summarize the consequences of this model as follows.  a) For the typical mass ratio the dominance of $L$ over $S_{1}$ is reversed as the separation in the binary decreases throughout the inspiral. In the last stages of the inspiral the spin dominates over the orbital angular momentum $S_{1}\\gg L$.  b) The angle $\\alpha $ between the orbital angular momentum and total angular momentum $\\mathbf{J}$ (with magnitude $J$), also the angle $\\beta $ between the dominant spin and total angular momentum evolve as:% \\begin{eqnarray} \\dot{\\alpha} &=&-\\frac{\\dot{L}}{J}\\sin \\alpha >0~,  \\label{alphadot} \\\\ \\dot{\\beta} &=&\\frac{\\dot{L}}{J}\\sin \\alpha <0~.  \\label{betadot} \\end{eqnarray}  c) The approximate expression relating $\\alpha $ to the post-Newtonian parameter $\\varepsilon $, mass ratio $\\nu $ and initial angle $\\alpha +\\beta $ span by the dominant spin with the orbital angular momentum (this angle being a constant during the inspiral) is:% \\begin{equation} \\tan \\alpha \\approx \\frac{\\sin \\left( \\alpha +\\beta \\right) }{\\varepsilon ^{-1/2}\\nu +\\cos \\left( \\alpha +\\beta \\right) }~.  \\label{alpha} \\end{equation}% (In Eq. (41) of Ref. \\cite{SpinFlip} the left hand side was given as $\\sin 2\\alpha /\\left( 1+\\cos 2\\alpha \\right) $.)  In a criticism to the work presented in Ref. \\cite{SpinFlip}, Gopakumar recently argued that \"it is unlikely that the spin-flip phenomenon will occur during the binary black hole inspiral phase\" \\cite{GopakumarConf}. This misconception comes from mixing up the \\textit{instantaneous} change in the direction of the total angular momentum, $\\,d\\mathbf{\\hat{J}}/dt=\\left( \\dot{L}/J\\right) \\left[ \\mathbf{\\hat{L}}-\\left( \\mathbf{\\hat{L}}\\cdot \\mathbf{\\hat{J}}\\right) \\mathbf{\\hat{J}}\\right] \\neq 0$ with its \\textit{% averaged} expression $\\left\\langle d\\mathbf{\\hat{J}}/dt\\right\\rangle =0$ over the precessional time-scale. The angles $\\alpha $ and $\\beta $ are \\textit{not} constants during the post-Newtonian evolution, as claimed in Ref. \\cite{GopakumarConf}, they rather change as given in Eqs. (\\ref% {alphadot})-(\\ref{betadot}).\\footnote{% Only when the total and orbital angular momenta are aligned, become the angles $\\alpha $ and $\\beta $ individually constant, as they identically vanish. Therefore in the aligned configuration no spin-flip could ever occur by the combined mechanism of SO precession and gravitational radiation.}  In the present paper we revisit some of the arguments of the spin-flip mechanism and also provide more details on it as compared to Ref \\cite% {SpinFlip}. In Section 2 we revisit the typical mass ratio argument, following a recent statistics of supermassive black hole candidates, resulting a newly established mass distribution. We comment on how these findings would affect the typical mass ratio range at SMBH encounters. In Section 3 we analyze how the spin-flip angle depends on the mass ratio and relative orientation of the spin and orbital angular momentum. We also derive a formula for the final spin during the inspiral. Finally we present our Concluding Remarks. ",
        "conclusions": "In light of the new data on a large sample of SMBH candidates we have established that the mass ratios obey an even logarithmic distribution in $% \\nu $. In the mass ratio range $\\nu \\in \\left( 1/30,1/3\\right) $ of SMBH mergers representing $40\\%$ of all possible cases, we have investigated the SO precession driven conservative and gravitational radiation driven dissipative contributions to the orbital evolution during the inspiral, averaged over the precession time-scale. In this mass range the ratio of the dominant spin magnitude and orbital angular momentum magnitude $S/L$ changes from \\textit{less than 1} to \\textit{larger than 1} during the inspiral. As the direction of the total angular momentum is unchanged on all time-scales larger than the precession time-scale, while the magnitude of the the orbital angular momentum decreases due to gravitational radiation and the magnitude of the spin stays constant, the spin direction has to change. The spin-flip of the dominant spin therefore occurs during the \\textit{inspiral}% . If jet activity is involved, X-shaped radio galaxies arise by this mechanism and a large set of observations on X-shaped radio galaxies could be explained.  In another $40\\%\\,$ of the mergers with mass ratios $\\nu \\in \\left( 1/1000,1/30\\right) $ the spin-flip never occurs by this mechanism, while in the remaining $20\\%$ of mergers with mass ratios $\\nu \\in \\left( 1/3,1\\right) $ it may occur during the plunge.  SMBH mergers of equal mass to $\\nu=1/3$ are only half as likely as the mass ratios $1/30$ to $1/3$, therefore the occurrence of the spin-flip can be considered typical during the inspiral. We analyzed the magnitude of the spin-flip angle occurring during the inspiral as function of the mass ratio and original relative orientation of the spin and orbital angular momentum and supported by numerical analysis the theoretical prediction (Fig \\ref% {Fig2}). We also derived a formula for the \\textit{final spin }at the end of the inspiral in this mass ratio range.  During the inspiral the following relations among the relevant time-scales hold: tilt / spin-flip time-scale $\\ge $ inspiral time-scale $\\gg $ precession time-scale $\\gg $ orbital time-scale (for all mass ratios in the typical range). Interestingly enough, the spin-flip time-scale for a typical mass ratio of $1/10$ is only about three years, while the precession time-scale is less then a day \\cite{SpinFlip}. Thus rapidly rotating relativistic jets coming close to our line of sight could produce significant variability at all wavelengths years before the coalescence. Therefore electromagnetic counterparts / precursors to the strongest gravitational wave emission are also likely to occur."
    },
    "1005/1005.3281_arXiv.txt": {
        "abstract": "On 2008 May 2, {\\em Chandra} observed the X-ray spectrum of $\\xi$~Boo (G8~V+K4~V), resolving the binary for the first time in X-rays and allowing the coronae of the two stars to be studied separately.  With the contributions of $\\xi$~Boo~A and B to the system's total X-ray emission now observationally established (88.5\\% and 11.5\\%, respectively), consideration of mass loss measurements for GK dwarfs of various activity levels (including one for $\\xi$~Boo) leads to the surprising conclusion that $\\xi$~Boo~B may dominate the wind from the binary, with $\\xi$~Boo~A's wind being very weak despite its active corona.  Emission measure distributions and coronal abundances are computed for both stars and compared with {\\em Chandra} measurements of other moderately active stars with G8-K5 spectral types, all of which exhibit a narrow peak in emission measure near $\\log T=6.6$, indicating that the coronal heating process in these stars has a strong preference for this temperature.  As is the case for the Sun and many other stars, our sample of stars shows coronal abundance anomalies dependent on the first ionization potential (FIP) of the element.  We see no dependence of the degree of ``FIP effect'' on activity, but there is a dependence on spectral type, a correlation that becomes more convincing when moderately active main sequence stars with a broader range of spectral types are considered.  This clear dependence of coronal abundances on spectral type weakens if the stellar sample is allowed to be contaminated by evolved stars, interacting binaries, or extremely active stars with $\\log L_X>29$, explaining why this correlation has not been recognized in the past. ",
        "introduction": "At a distance of only 6.7~pc, the primary of the $\\xi$~Boo binary system, $\\xi$~Boo~A (G8~V), is one of the two brightest ``moderately active'' coronal X-ray sources in the sky.  The only star of this coronal activity level with a comparable X-ray flux at Earth is $\\epsilon$~Eri (K2~V, $d=3.2$~pc).  As such, these two stars are of crucial importance for establishing the coronal properties of stars that are significantly more active than the Sun, but which are not in the extremely active regime that would be represented by the RS~CVn binaries or very young stars with rotation periods of only a few days.  \\citet{jjd00} find that the coronal emission measure (EM) distributions of $\\epsilon$~Eri and $\\xi$~Boo are similar to those of bright active regions on the Sun.  They note that the X-ray emission from these stars can therefore be explained by stellar surfaces that are completely covered with solar-like active regions, suggesting that $\\epsilon$~Eri and $\\xi$~Boo~A represent a maximum of solar-like activity for coronal stars, whereas stars with even higher X-ray luminosities and higher coronal temperatures represent a completely different regime of coronal activity.  This interpretation places $\\epsilon$~Eri and $\\xi$~Boo~A at an interesting coronal transition point.  Further evidence that $\\epsilon$~Eri and $\\xi$~Boo~A lie near a transition in the nature of coronal activity comes from stellar wind measurements, which are based on analyses of H~I Ly$\\alpha$ absorption from wind-ISM interaction regions \\citep{bew05a}.  At low activity levels mass loss increases with activity, with $\\epsilon$~Eri having a mass loss rate 30 times that of the Sun (i.e., $\\dot{M}=30~\\dot{M}_{\\odot}$, where $\\dot{M}_{\\odot}=2\\times 10^{-14}$ M$_{\\odot}$~yr$^{-1}$).  However, the somewhat more active $\\xi$~Boo system has a surprisingly weak wind with only $5~\\dot{M}_{\\odot}$, leading to the suggestion that the observed increase of mass loss with activity ends abruptly at an X-ray surface flux of about $F_X=1\\times 10^{6}$ ergs cm$^{-2}$ s$^{-1}$, with $\\epsilon$~Eri and $\\xi$~Boo on opposite sides of this dividing line \\citep{bew05a}.  The wind measurements suggest a slight modification of the proposal that $\\epsilon$~Eri and $\\xi$~Boo~A represent an extreme of truly solar-like activity.  They suggest that $\\epsilon$~Eri is the true extreme, while $\\xi$~Boo~A actually lies beyond it at the low end of the non-solar-like activity regime.  Evidence that $\\xi$~Boo~A's magnetic topology is fundamentally different from the Sun comes from spectropolarimetric measurements of photospheric fields on $\\xi$~Boo~A.  \\citet{pp05} find evidence for both a global dipole field component much stronger than that of the Sun ($\\sim 40$~G), and a strong toroidal field component ($\\sim 120$~G) that has no clear solar analog.  We also note that the activity level of $\\xi$~Boo~A is roughly where one starts to see polar starspots \\citep{kgs02}, which also have no solar analog.  \\citet{cgt88} find evidence for high latitude spots on $\\xi$~Boo~A, though not necessarily ``polar spots.''  Perhaps the strong large-scale fields found by \\citet{pp05} envelope the star and inhibit wind flow, explaining why $\\xi$~Boo's wind is surprisingly weak.  Consistent with this idea, \\citet{pp08} present evidence that as stellar rotation and magnetic activity increase, stars tend to store a larger fraction of their magnetic energy in large scale fields, which could in principle inhibit mass loss.  With this background in mind, we here present an analysis of the X-ray spectrum of $\\xi$~Boo~A based on new observations from {\\em Chandra}, with the goal of seeing whether the apparent changes in coronal character noted above find clear expression in the coronal spectrum of $\\xi$~Boo~A, and coronal properties inferred from it.  A similar spectrum of $\\epsilon$~Eri has already been analyzed, so a central goal of this paper is to compare the $\\xi$~Boo~A spectrum with that of $\\epsilon$~Eri and similar stars \\citep[][hereafter WL06]{bew06}.  Other coronal spectra of $\\xi$~Boo have been obtained by previous X-ray and EUV missions \\citep[e.g.,][]{jml99,jcp09}, but the new {\\em Chandra} data are superior in many ways.  One of them is that {\\em Chandra}/LETGS observations allow the X-ray spectrum to be observed at high spectral resolution over a very broad spectral range of $5-175$~\\AA. This allows the detection of many coronal emission lines that have never been previously detected before from $\\xi$~Boo~A.  Another particulary important improvement is that {\\em Chandra}'s superb spatial resolution allows us for the first time to truly address the issue of $\\xi$~Boo's binarity.  For $\\xi$~Boo~A has a companion, $\\xi$~Boo~B (K4~V), which is only 6$^{\\prime\\prime}$ away, and {\\em no X-ray or UV observation of $\\xi$~Boo has ever truly resolved the binary.}  Table~1 lists various properties of the two $\\xi$~Boo stars.  The faster rotation and larger size of the primary lead to the expectation that $\\xi$~Boo~A should dominate the coronal emission from the system.  Detailed analyses of IUE data and a ROSAT/HRI image of $\\xi$~Boo do indeed suggest that $\\xi$~Boo~A dominates the binary's UV and X-ray emission \\citep{lh79,js97}, so the usual assumption of simply assigning all flux to the primary should be a decent approximation.  However, $\\xi$~Boo~B's spectral type and rotation period are not that different from the X-ray luminous $\\epsilon$~Eri (K2~V; $P_{rot}=11.7$~days).  Thus, $\\xi$~Boo~B should be contributing a non-negligible flux to the system's total X-ray emission.  The older observations simply do not have sufficient spatial resolution to quantify how much.  In short, the new {\\em Chandra} data presented here allow the first X-ray detection of $\\xi$~Boo~B, and they ensure that the brighter spectrum of $\\xi$~Boo~A is uncontaminated by any emission from the companion. ",
        "conclusions": "We have used {\\em Chandra}/LETGS observations to study the coronae of the $\\xi$~Boo binary, and we have compared our results with a previous study of five other moderately active GK dwarfs: $\\epsilon$~Eri, 70~Oph~AB, and 36~Oph~AB.  Our findings are summarized as follows: \\begin{enumerate} \\item The two stellar components of $\\xi$~Boo are resolved for the first time in X-rays.  We find that $\\xi$~Boo~A accounts for 88.5\\% of the counts in the zeroth order image.  For comparison, spectral analysis implies an 84.6\\% contribution to the system's X-ray flux in the ROSAT-like 0.1-2.4~keV bandpass. \\item With $\\xi$~Boo~B detected for the first time in X-rays, we now know it has an X-ray luminosity comparable to stars known to have strong stellar winds, leading to the conclusion that $\\xi$~Boo~B may account for most of the modest wind observed from the system, with the more active $\\xi$~Boo~A possessing a surprisingly weak wind despite its active corona.  This strengthens the case for the presence of a ``wind dividing line'' separating moderately active stars with strong winds and very active stars with weak winds, possibly due to the increasing strength of large scale global magnetic fields that inhibit mass loss. \\item Our emission measure analysis of $\\xi$~Boo~A and B finds evidence for a narrow peak in EM at $\\log T=6.6$ for both stars, a peak also seen for the other moderately active G8-K5 dwarfs in our sample, indicating that the coronal heating mechanism operating on these stars favors this temperature. \\item The emission measure analysis provides us with coronal abundances relative to Fe.  Looking for coronal abundance anomalies requires correcting for stellar photospheric abundances, which in turn involves assuming solar photospheric abundances, since the stellar abundances we use are measured relative to solar.  When we simply assume solar abundances from \\citet{ng98}, for our sample of stars we find that the high-FIP element abundances are inconsistent with each other.  This inconsistency is decreased significantly for all our stars if we instead assume a photospheric Ne abundance consistent with stellar coronal measurements from \\citet{jjd05}, and solar abundances from \\citet{ma09} for other elements.  This is an argument in favor of these latter photospheric abundances, unless there are actual coronal fractionation processes taking place among the high-FIP elements that we are erroneously removing. \\item Among low-FIP ions, the coronal Mg/Si ratio is consistently higher than the photospheric ratio for all of our G8-K5 stars, presumably an effect of actual coronal fractionation.  An enhancement of Mg relative to Si in the corona is predicted by many of the models of \\citet{jml09}, which are designed to explain solar/stellar coronal abundance anomalies.  However, although \\citet{at05} also finds high Mg/Si ratios for early G dwarfs, coronal Mg/Si ratios appear to be {\\em lower} than photospheric for M dwarfs \\citep{cl08}, suggesting a spectral type dependence for this coronal abundance ratio. \\item With regards to the FIP effect, our sample of stars exhibit a range of FIP bias, with 70~Oph~A and $\\xi$~Boo~A having the strongest solar-like FIP effect, $\\epsilon$~Eri showing little or no FIP effect, and 70~Oph~B possessing a slight inverse FIP effect. Expanding our sample of stars even further to include early G stars from \\citet{at05} and M dwarfs from \\citet{cl08}, we find that the FIP bias correlates beautifully with spectral type, but only if extremely active stars with $\\log L_X>29$ are excluded. \\item The tight correlation of FIP bias with spectral type suggests to us that the nature and causes of the FIP effect can ultimately be traced directly to fundamental surface properties.  Among normal main sequence stars, these change in a very predictable manner with spectral type, explaining the tight correlation in Figure~9.  No such correlation exists for extremely active stars, because many are evolved stars with disparate photospheric characteristics, and in other cases the stellar atmospheres will be affected by rapid rotation or the effects of close binarity. \\end{enumerate}"
    },
    "1005/1005.4232_arXiv.txt": {
        "abstract": "A number of new and planned radio telescopes will consist of large arrays of low-gain antennas operating at frequencies below 300~MHz.  In this frequency regime, Galactic noise can be a significant or dominant contribution to the total noise.  This, combined with mutual coupling between antennas, makes it difficult to predict the sensitivity of these instruments.  This paper describes a system model and procedure for estimating the system equivalent flux density (SEFD) -- a useful and meaningful metric of the sensitivity of a radio telescope -- that accounts for these issues.  The method is applied to LWA-1, the first ``station'' of the Long Wavelength Array (LWA) interferometer.  LWA-1 consists of 512 bowtie-type antennas within a 110 $\\times$ 100~m elliptical footprint, and is designed to operate between 10~MHz and 88~MHz using receivers having noise temperature of about 250~K.  It is shown that the correlation of Galactic noise between antennas significantly desensitizes the array for beam pointings which are not close to the zenith.  It is also shown that considerable improvement is possible using beamforming coefficients which are designed to optimize signal-to-noise ratio under these conditions.  Mutual coupling is found to play a significant role, but does not have a consistently positive or negative influence.  In particular, we demonstrate that pattern multiplication (assuming the behavior of single antennas embedded in the array is the same as those same antennas by themselves) does not generate reliable estimates of SEFD. ",
        "introduction": "Introduction}  A number of new and planned radio telescopes will consist of large arrays of closely-spaced low-gain antennas operating at frequencies below 300~MHz.  These include LWA \\cite{PIEEE_LWA}, LOFAR \\cite{PIEEE_LOFAR},  MWA \\cite{PIEEE_MWA}, and SKA \\cite{PIEEE_SKA}.  In this frequency regime, it is possible to design receivers with noise temperatures that are much less than the antenna temperatures associated with the ubiquitous Galactic synchrotron radiation, such that the resulting total system noise temperature is dominated by Galactic noise \\cite{E05}.  This is quite different from the condition most often considered, in which it is usually assumed that internal noise associated with the receivers dominates and does not scatter into the array, so that the noise associated with different antennas is uncorrelated; see e.g. \\cite{VB88,Lee93,Kraft00}.  Recent work on arrays for long-wavelength radio astronomy accounts for the dominance of external noise, but neglects the effects of correlation of this noise between antennas \\cite{VBA05,PIEEE_LWA}.  However, it is known that this correlation is likely to have a significant effect for these instruments \\cite{LWA142}, and this issue is further explored in this paper.  In \\cite{CPD04,C05}, and \\cite{IMW08}, correlation of noise between antennas is considered, but the source of noise is internal (amplifiers or ohmic losses in antennas) and the correlation arises due to propagation internal to the array.  The problem of correlation of external noise has been intermittently considered in communications (e.g. \\cite{Gans06}) and direction finding (e.g. \\cite{AST07}) applications, but does not seem to have been previously considered for long-wavelength radio astronomy beamforming applications.  Because antenna spacing in the systems of interest is typically less than a few wavelengths, mutual coupling also plays a significant role.  Since these arrays are electromagnetically large, interact with the electromagnetically-complex ground, and may have aperiodic spacings, it is difficult to determine the characteristics of antennas, either individually or collectively as part of a beamforming system. In particular, it is usually difficult to know if pattern multiplication -- that is, assuming that the behavior of single antennas embedded in the array is the same as those same antennas by themselves -- yields reasonable results.  Past studies have shown that mutual coupling in aperiodic arrays of low-gain elements results in fluctuation of beam gain and sidelobe levels as a function of scan angle when element spacing is less than a few wavelengths \\cite{LS67,AL72}.  This suggests that pattern multiplication may not be a useful assumption. However, useful and generalizable findings which are applicable to the systems of interest are not commonly available.  This paper describes a procedure for estimating the sensitivity of radio telescope arrays which is appropriate under these conditions.  The procedure is based on a system model, described in Section~\\ref{sTheory}, which relates the electromagnetic response of the array (the {\\it array manifold}), a model for the external noise temperature, and a model for the receiver noise temperatures to the system equivalent flux density (SEFD) achieved by a beam formed using specified beamforming coefficients. SEFD is defined as the power flux spectral density (i.e., W~m$^{-2}$~Hz$^{-1}$) which yields signal-to-noise ratio (SNR) equal to unity at the beamformer output.  SEFD is a useful metric as it includes the combined effect of antennas and all noise sources into a single ``bottom line'' number that is directly related to the sensitivity of astronomical observations.  The primary difficulty in determining SEFD using the above procedure is obtaining the array manifold.  One approach is demonstrated by example in Section~\\ref{sLWA} of this paper.  We analyze LWA-1, the first ``station'' of the Long Wavelength Array (LWA) interferometer \\cite{PIEEE_LWA}.  LWA-1 consists of 512 bowtie-type antenna elements arranged into 256 dual-polarized ``stands'' within a 110 $\\times$ 100~m elliptical footprint.  LWA-1 is designed to operate between 10~MHz and 88~MHz using receivers having noise temperature of about 250~K.  We obtain the array manifold for LWA-1 at 20~MHz, 38~MHz, and 74~MHz using a method of moments (MoM) wire-grid model.  Because the model is too large to analyze all at once (a common problem with this class of arrays), a procedure described in Section~\\ref{ssArray} is employed in which the manifold is calculated for one stand (i.e., one pair of collocated antenna elements) at a time.  Unlike pattern multiplication in which the presence of the remaining stands would be ignored, this procedure obtains the response of each antenna in the presence of nearby antennas and structures.  Also considered in this paper is the selection and performance of beamforming coefficients which optimize SEFD.  Because mutual coupling and external noise correlation are significant, it is to be expected that ``simple'' beamforming coefficients based solely on antenna positions (i.e., phases associated with geometry only) will not be optimal.  In Section~\\ref{sAP}, the SEFD performance of LWA-1 is evaluated.  It is shown that the optimal coefficients significantly improve sensitivity relative to simple coefficients.  Finally, in Section~\\ref{sConc}, the extension of these results to predict the imaging performance of an interferometer comprised of multiple beamforming arrays is considered. ",
        "conclusions": ""
    },
    "1005/1005.3909_arXiv.txt": {
        "abstract": "We present a stellar population synthesis study of a type I luminous infrared galaxy (LIRG): IRAS\\,F13308+5946. It is a quasar with absolute magnitude $M_i=-22.56$ and has a spectral feature of a Seyfert 1.5 galaxy. Optical images show characteristics of later stages of a merger. With the help of the stellar synthesis code \\starlight\\ (Cid Fernandes et al.\\ 2005) and both Calzetti et al.\\ (2000) and Leitherer et al.'s (2002) extinction curves, we estimate the past infrared (IR) luminosities of the host galaxy and find it may have experienced an ultraluminous infrared galaxy (ULIRG) phase for nearly 300\\,Myr, so this galaxy has probably experienced a type I ULIRG phase. Both nuclear starburst and active galactic nuclei (AGN) contribute to the present IR luminosity budget, and starburst contributes $\\sim$70\\%. The mass of supermassive black-hole (SMBH) is $M_{\\rm BH}=1.8\\times10^8\\,M_\\odot$ and the Eddington ratio $L_{\\rm bol}/L_{\\rm Edd}$ is 0.12, which both approximate to typical values of PG QSOs. These results indicate that IRAS\\,F13308+5946 is probably at the transitional phase from a type I ULIRG to a classical QSO. ",
        "introduction": "Luminous infrared galaxies (LIRGs) and ultraluminous infrared galaxies (ULIRGs) are galaxies with infrared (IR) luminosities $L_{\\rm IR} = L(8\\text{--}1000\\mu{\\rm m}) > 10^{11}$\\,$L_{\\odot}$ and $L_{\\rm IR} > 10^{12}$\\,$L_{\\odot}$, respectively. Their IR luminosities and space densities are both comparable to QSOs (Sanders \\& Mirabel 1996). These galaxies always show signs of tidal interaction and merger, and the interaction rate is observed to increase with IR luminosity (Zou et al.\\ 1991; Clements et al.\\ 1996; Murphy et al.\\ 1996; Wu et al.\\ 1998; Veilleux et al.\\ 2002). Nearly all ULIRGs are strong interaction or merger systems (Kim et al.\\ 2002).  It is widely accepted that galaxy interactions and mergers trigger extreme nuclear activity, as well as more widespread starburst (Toomre \\& Toomre 1972; Larson \\& Tinsley 1978). The molecular gas concentrations of ULIRGs are in their central kpc regions (Downes \\& Solomon 1998) and have the ability to form stars with densities comparable to those in elliptical galaxies (Tacconi et al.\\ 2002). Kormendy \\& Sanders (1992) proposed that ULIRGs may evolve into elliptical galaxies through merger induced, dissipative collapse. Structural, kinematic and photometric properties of ULIRGs show that they are originated from major mergers and they averagely fall on the fundamental plane of moderate-mass ellipticals (stellar mass $\\sim 10^{10}$--$10^{11}\\,M_\\odot$), but are well offset from giant ellipticals, suggesting that ULIRG mergers are ellipticals in formation (Genzel et al.\\ 2001; Tacconi et al.\\ 2002; Dasyra et al.\\ 2006a).  Observations revealed that a large fraction of LIRGs/ULIRGs show spectral characteristics as Seyfert galaxies (Wu et al.\\ 1998; Kewley et al.\\ 2001). The fraction increases dramatically with IR luminosity, and among ULIRGs, the fraction is about 50\\% (Kim et al.\\ 1998a; Veilleux et al.\\ 1999; Cao et al.\\ 2006; Yuan et al.\\ 2010; Nardini et al.\\ 2010), or even higher ($\\sim$70\\%; Nardini et al.\\ 2010). Though many of these sources are dominated by AGN in their bolometric luminosity (Boller et al.\\ 2002; Nandra \\& Iwasawa 2007), in most cases the predominance of starburst over AGN is proposed (Gu et al.\\ 1997; Lutz et al.\\ 1999; Franceschini et al.\\ 2003), and Nardini et al.\\ (2008) proposed a fraction of $\\sim85$\\%.  From millimeter-wave CO observations and optical spectra analyse of 10 ULIRGs, Sanders et al.\\ (1988) proposed that ULIRGs represent the dust-enshrouded stage of QSOs. They proposed a classical evolution scenario that two gas-rich spirals merge (or interact) first and the funneled gas toward the merger center triggers nuclear starburst before the ignition of a dust-enshrouded AGN. When dust has been consumed or swept away under the radiative pressure of AGN and supernovae, an optical quasar would appear. Canalizo \\& Stockton (2000, 2001) selected a sample of 9 low-redshift QSOs which fall onto the intermediate position between the regions occupied by ULIRGs and QSOs in the far-infrared (FIR) color-color diagram. All these 9 transition QSOs are undergoing tidal interactions and 8 are major mergers. All of them also show strong recent star formation activity within 300\\,Myr. They proposed a model involving a dust cocoon that initially surrounds the QSO nuclear regions, which can account for all the observed and derived properties in transition objects, and it is consistent with the idea of Sanders et al.\\ (1988). They suggested that either at least some ULIRGs evolve to become classical QSOs, or some QSOs are born under the conditions as ULIRGs and their lifetimes as QSOs last $\\leq $ 300\\,Myr. Most of these transition QSOs are also ULIRGs. As Kawakatu et al.\\ (2006), we refer to QSOs and Seyfert 1 galaxies that selected from ULIRGs as type I ULIRGs (or IR QSOs in Zheng et al.\\ 2002). Kawakatu et al.\\ (2006) and Hou et al.\\ (2009) find that type I ULIRGs are the early phase of black-hole (BH) growth and they are QSOs in formation.  The case may be as Colina et al.\\ (2001) had proposed that high-luminosity QSOs would be the end point in the merging process of massive ($>L$*) disk galaxies, while cool ULIRGs, which would be the end product in the merging of two or more low-mass (0.3$L$*--0.5$L$*) disk galaxies, would not evolve into QSOs. This scenario is supported by many studies (McLeod \\& Rieke 1994; Arribas et al.\\ 2003; Dasyra et al.\\ 2006a).  In this paper we present the result of stellar population synthesis of IRAS\\,F13308+5946 and estimations of past IR luminosities. IRAS\\,F13308+5946 is a LIRG at present and shows a spectral characteristic of a Seyfert 1.5 galaxy, which together suggest a potential transition phase between type I ULIRGs and classical QSOs. It shows clear absorption lines and young stellar population features for spectral fitting. Observations from {\\it IRAS}, 2MASS, SDSS and FIRST are presented in \\S2. We describe the process of extracting the star formation history (SFH), emission-line fitting, aperture correction and the extinction curve in \\S3. In \\S4 we present some estimations like the current star formation rate (SFR), the BH mass, the Eddington ratio, etc., based on the emission-lines fitting. The past IR luminosities are estimated in \\S5. Conclusions and discussions are given in \\S6. We summarize in \\S7. Through out this paper we adopt a cosmology of $H_0$=71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$=0.29 and $\\Omega_{\\Lambda}$=0.71. ",
        "conclusions": "\\subsection{Fe $_{\\rm II}$ pseudo-continuum} Many IR QSOs are extremely strong Fe $_{\\rm II}$ emitters (Zheng et al.\\ 2002; L\\'ipari et al.\\ 2003). In IRAS\\,F13308+5946 spectrum, we also find Fe $_{\\rm II}$ pseudo-continuum emission in wavelength region 4270--4700\\AA\\ and 5100--5600\\AA. We have masked these two wavelength regions before spectral synthesis. Actually, we have attempted to subtract Fe $_{\\rm II}$ lines from $O_\\lambda$ before fitting. However, as the case of IRAS\\,Z11598$-$0112, IRAS\\,F02065+4705, etc.\\ (Zheng et al.\\ 2002), Fe $_{\\rm II}$ multiplets 37, 38 (4500--4680\\AA) are relatively stronger than the Fe $_{\\rm II}$ multiplets 48, 49 (5100--5400\\AA) compared with Boroson \\& Green (1992, hereafter BG92) Fe $_{\\rm II}$ template, and also the V\\'eron-Cetty et al.\\ (2004) template. Thus, the Fe $_{\\rm II}$ multiplets 37, 38 are left when the multiplets 48, 49 are removed. We have also attempted to remove Fe $_{\\rm II}$ lines from the residual spectrum, but the result was not improved. Furthermore, including Fe $_{\\rm II}$ template as a component to fit $O_\\lambda$ does not bring any improvement at all. Therefore, we masked the Fe $_{\\rm II}$ pseudo-continuum regions before spectral synthesis, as have mentioned in \\S3.4.  The Fe $_{\\rm II}$ $\\lambda$4570 (4434--4684\\AA) flux is measured during the procedures above and the line ratio Fe $_{\\rm II}~\\lambda$4570/H$\\beta$ is 0.81. This ratio is typical of optical selected QSOs (BG92). We find that the Fe $_{\\rm II}$ continuum extension to [O $_{\\rm II}$] $\\lambda3727$, H$\\beta$, etc.\\ lines is hard to detect, so we measure these lines fluxes directly.  \\subsection{Extinction Curve} In this paper we don't adopt Cardelli et al.'s (1989) extinction curve, because it is for cases which are similar to Milky Way and not applicable to starburst galaxies. We use the Calzetti curve and Leitherer et al.'s (2002) extension form, because they are derived from star-forming regions and starburst galaxies.  Dust in different components and stellar populations vary in temperature, ingredient, amount, geometry, etc., so different attenuation for each component should be adopted if we could distinguish them. However, as the spectra from the host galaxy and central AGN are coupled, and also there is only one extinction term in \\starlight, the extinction obtained from spectral synthesis is a combination value from both regions. Also, there may exist some optically thick star formation regions, which may not be detected even by optical or near-infrared, but they still contribute to IR flux. In this case, the actual extinction could be underestimated.  From Balmer decrement, the line ratio H$\\alpha/$H$\\beta$, the derived color excess of stellar continuum  is $\\esbv = 0.3115$, which is twice as the one obtained from spectral fitting. This indicates differential attenuation between the dusty starburst region and older populations.  \\subsection{IRAS\\,F13308+5946: A Possible Evolutionary Link between (type I) ULIRGs and QSOs} The spectral synthesis gives a 500\\,Myr starburst history, which implies a triggering mechanism took place during such epoch. Galaxy interactions and mergers have timescales of a few $\\times 10^8$\\,yr (Binney \\& Tremaine 1987). The similar timescales imply that the starburst was possibly triggered by galaxies merger.  In \\S4 and \\S5, we demonstrated that IRAS\\,F13308+5946 has possibly experienced an IR QSO phase since $\\sim$300\\,Myr ago. Follow Hao05, we plot IRAS\\,F13308+5946 on the diagram which shows the relation between IR luminosity and the bolometric luminosity measurement (represented by $\\lambda L_\\lambda(5100\\AA)$; Kaspi et al.\\ 2000) in Fig.~\\ref{f9eps}. IRAS\\,F13308+5946 locates at a transitional position between IR QSOs and PG QSOs. The tight correlation followed by PG QSOs and NLS1s suggests their IR luminosities are associated with the optical through central AGNs, while IR QSOs have IR excess powered by other mechanisms rather than AGN. As for IRAS F13308+5946, we attribute the IR excess to the ongoing star formation. In Fig.~\\ref{f10eps}, which shows the IR spectral index $\\alpha(60, 25)$ (defined as $\\alpha(\\lambda_1, \\lambda_2)=-{{\\rm log}[F(\\lambda_2)/F(\\lambda_1)] \\over {\\rm log}(\\lambda_2/\\lambda_1)}$) versus the IR excess $L_{\\rm IR}/L_{\\rm 5000}$, it also locates between IR QSOs and PG QSOs. The IR excess increase as the dust temperature decrease, suggesting that star formation activities become more and more important in powering IR luminosity of IR QSOs and ULIRGs, because the temperature of dust heated by stars is lower than that by AGN.  On the other hand, Hao05 gives statistical median $M_{\\rm BH}$ values of IR QSOs and PG QSOs, which are 4.9$\\times10^7\\,M_\\odot$ and 2.1$\\times10^8\\,M_\\odot$. Thus, IRAS\\,F13308+5946 has both BH mass and Eddington ratio similar to those of PG QSOs. Furthermore, as a type 1.5 galaxy, the featureless continuum and emission-lines from the AGN may be partially attenuated by the dusty torus, and the AGN itself may be powering a classical QSO, for the galaxy is an {\\it i}-band quasar.  Therefore, IRAS\\,F13308+5946 may evolve into an optical QSO when the starburst ceases and the nuclear dust is dissipated by the radiative pressure of AGN. In Fig.~\\ref{f9eps} and \\ref{f10eps}, the positional arrangement of ULIRGs, IR QSOs and classical QSOs implies an evolutionary sequence: ULIRGs $\\rightarrow$ type I ULIRGs / IR QSOs $\\rightarrow$ (IRAS\\,F13308+5946) $\\rightarrow$ PG QSOs. As the {\\it H}-band luminosity is $\\sim$4$L$*, it support the evolutionary scenario proposed by Colina et al.\\ (2001) that ULIRGs generated by two or more massive ($\\geq L_*$) galaxies would evolve into QSOs (also supported by Sanders et al.\\ 1988 and Lutz et al.\\ 1999). \\begin{figure}[h] \\includegraphics[width = 84mm]{f9.eps} \\caption{IR luminosity vs. the bolometric luminosity measurement, $\\lambda L_\\lambda$5100\\AA. The {\\it red} open circles represent IR QSOs; the {\\it green} filled circles represent PG QSOs; the {\\it brown} filled triangles represent NLS1s; the solid line represents the linear regression for all PG QSOs and NLS1s. These data points are reproduced with the data from Hao05. IRAS\\,F13308+5946 is represented by the {\\it blue} star.} \\label{f9eps} \\end{figure} \\begin{figure}[h] \\includegraphics[width = 84mm]{f10.eps} \\caption{IR spectral index $\\alpha(60, 25)$ vs. the IR excess, $L_{\\rm IR}/L_{\\rm 5100}$. The {\\it purple} open triangles represent 15 ULIRGs, which are selected from Kim et al.\\ (1998b) and have SDSS spectra; other data points are also from Hao05 and have the same meaning as Fig.~\\ref{f9eps}; the three {\\it red} open circles enclosed by {\\it green} diamonds represent IR QSOs which are also PG QSOs. IRAS\\,F13308+5946 is represented by the {\\it blue} star and marked with a lower limit.} \\label{f10eps} \\end{figure}  Finally, through {\\it r}-band surface photometry, we find the outskirt surface brightness profile follows the de Vaucouleurs law and the central region follows a point spread function (PSF). It is consistent with other studies which show that some (maybe even higher fraction) ULIRGs may undergo QSOs phase in their evolutionary history before they settle down as ellipticals (Zheng et al.\\ 1999; Arribas et al.\\ 2004; Dasyra et al 2006b)."
    },
    "1005/1005.2215_arXiv.txt": {
        "abstract": "{ We calculate the relic abundance of thermally produced neutralino cold dark matter in the general 19 parameter supergravity (SUGRA-19) model. A scan over GUT scale parameters reveals that models with a bino-like neutralino typically give rise to a dark matter density $\\Omega_{\\tz_1}h^2\\sim 1-1000$, {\\it i.e.} between 1 and 4 orders of magnitude higher than the measured value. Models with higgsino or wino cold dark matter can yield the correct relic density, but mainly for neutralino masses around 700-1300 GeV. Models with mixed bino-wino or bino-higgsino CDM, or models with dominant co-annihilation or $A$-resonance annihilation can yield the correct abundance, but such cases are extremely hard to generate using a general scan over GUT scale parameters; this is indicative of high fine-tuning of the relic abundance in these cases. Requiring that $m_{\\tz_1}\\alt 500$ GeV (as a rough naturalness requirement) gives rise to a minimal probably dip in parameter space at the measured CDM abundance. For comparison, we also scan over mSUGRA space with four free parameters. Finally, we investigate the Peccei-Quinn augmented MSSM with mixed axion/axino cold dark matter. In this case, the relic abundance agrees more naturally with the measured value. In light of our cumulative results, we conclude that future axion searches should probe much more broadly in axion mass, and deeper into the axion coupling. } ",
        "introduction": "\\label{sec:intro}   The cosmic abundance of cold dark matter (CDM) has been recently extracted at high precision by the WMAP collaboration\\cite{wmap7}. An analysis of WMAP seven year data implies \\be \\Omega_{CDM}h^2=0.1123\\pm 0.0035\\ \\ \\ (68\\%\\ CL) , \\ee where $\\Omega=\\rho/\\rho_c$ is the dark matter density relative to the closure density, and $h$ is the scaled Hubble constant. Since no particle contained within the Standard model of particle physics has exactly the right properties to constitute CDM, new matter states from physics beyond the SM are required for explanation\\cite{dmreview}.  The most popular explanation of CDM comes from models including weak scale supersymmetry (SUSY)\\cite{wss}. SUSY models are invoked to cure the hierarchy problem contained within the SM. SUSY models also receive some indirect support from experiment\\cite{gauge} in that the gauge couplings, measured at the scale $Q=M_Z$, meet nearly at a point when run to $M_{GUT}$ via the renormalization group equations (RGEs) of the Minimal Supersymmetric Standard Model, or MSSM\\cite{drw}. Within $R$-parity conserving SUSY theories, there are several dark matter candidates, including the lightest neutralino $\\tz_1$, the gravitino $\\tG$ and the axion/axino admixture $a\\ta$ from the Peccei-Quinn\\cite{pq,ww} augmented MSSM (or PQMSSM).\\cite{steffen_review}  Most analyses focus upon the lightest neutralino $\\tz_1$ as a CDM candidate, since it can be classified as a weakly interacting massive particle, or WIMP. WIMP particles such as $\\tz_1$ are especially compelling due to the so-called ``WIMP miracle''. It is assumed the DM particle is present in thermal equilibrium at some early time in the universe's history, and then the relic abundance can be found by solving the Boltzmann equation for its number density. Particles with weak scale mass, which interact with weak scale strength, should give a relic abundance roughly in the vicinity of what is measured.  However, detailed calculations of the neutralino relic abundance in models such as mSUGRA\\cite{msugra} (or CMSSM) find that throughout parameter space, the neutralino abundance is usually much higher than the measured value, and typically by 1-2 orders of magnitude\\cite{inoDM}. Only in very narrow regions of parameter space is the measured CDM abundance found: the stau co-annihilation region\\cite{stau}, the hyperbolic branch/ focus point region with mixed higgsino DM\\cite{hb_fp}, the $A$-resonance region at large $\\tan\\beta$\\cite{Afunnel}, and perhaps at the light Higgs resonance\\cite{higgs}.  That the mSUGRA parameter regions where the predicted abundance agrees with measurement are quite narrow is indicative of a needed fine-tuning of fundamental model parameters to achieve the required value. The fine-tuning is a measure of the slope of the surface of $\\Omega_{\\tz_1}h^2$ as a function of model parameters $a_i$.  Quantitatively, the fine-tuning with respect to variation in the parameter $a_i$ is defined as \\begin{equation} \\Delta_{a_i} \\equiv \\frac{\\partial\\log\\Omega_{DM}h^2}{\\partial\\log a_i} . \\end{equation} The overall fine-tuning can be obtained by combining the $\\Delta_{a_i}$ in quadrature. The degree of fine-tuning of the relic abundance in mSUGRA has been calculated in Ref's \\cite{eo} and \\cite{bbox}. Both groups find that the regions of agreement between theory and experiment are highly fine-tuned. Ref. \\cite{bbox} claims also that the very large $\\tan\\beta$ ($\\agt 50$) region is not very fine-tuned, since the pseudoscalar Higgs width $\\Gamma_A$ becomes very large (owing to the large $b$-quark Yukawa coupling), and efficient neutralino annihilation can occur all over the $m_0\\ vs.\\ m_{1/2}$ plane.  In Ref. \\cite{bbs}, we proposed mixed axion/axino cold dark matter in the Peccei-Quinn augmented mSUGRA model. By requiring a re-heat temperature in the $10^6-10^8$ GeV range, a MeV-scale axino mass and a large value of $f_a\\agt 10^{11}$ GeV is favored. As opposed to the high fine-tuning which is present in mSUGRA for neutralino CDM, if instead mixed axion/axino CDM is invoked, with a $\\sim$ MeV scale axino\\cite{axmass}, then the fine-tuning needed to achieve a theory-experiment match is relatively low. The lowest fine-tuning is observed for the case where there is a nearly equal mix of axion and thermally produced axino DM\\cite{bbox}.  A major criticism of the result found in Ref. \\cite{bbox} is that the fine-tuning analysis of neutralino CDM is restricted to the mSUGRA model. It might be possible that if one opens up the parameter space to include models with non-universality, then many more possibilities may arise to gain a theory-experiment match in the dark matter relic density.  In this paper, we calculate the relic abundance of neutralino CDM in a 19 parameter version of the MSSM: the SUGRA-19 model. Another recent calculation of the neutralino CDM abundance in a 19 parameter MSSM has been performed by Berger {\\it et al.}\\cite{bghr}.\\footnote{ Similar analyses using weak scale SUSY parameters and requiring a thermal abundance of neutralino CDM in accord with WMAP are given in Ref's \\cite{py,aaqfh,bbps}.} In that analysis, the so-called pMSSM model is examined, which is a version of the MSSM with minimal $CP$ and flavor violating terms. The 19 soft SUSY breaking and other parameters are input at the weak scale, and it is required that $\\Omega_{\\tz_1}h^2\\le 0.11$. In the analysis presented here, we will also adopt a 19 parameter MSSM, but will insist on adopting soft SUSY breaking parameters {\\it at the GUT scale}. The reason is that the unification of gauge couplings provides circumstantial evidence that the MSSM, or MSSM plus gauge singlets (or additional $SU(5)$ multiplets) is the correct effective field theory between $M_{weak}$ and $M_{GUT}$. The pattern of soft term running as provided by the MSSM RGEs gives us some guidance as to how likely  are various sets of weak scale parameters. For instance, by sampling over weak scale parameters including the superpotential $\\mu$ term, one easily generates low $\\mu$ solutions which lead to higgsino-like neutralinos. However, sampling over GUT scale parameters, it is not easy to generate low $\\mu$ solutions with higgsino-like dark matter. The reason is that (at tree level) $\\mu^2(m_{weak})\\simeq -m_{H_u}^2(M_{weak})$, and it is rare that $m_{H_u}^2$ runs to just barely negative values, which would lead to a small $\\mu$ parameter. In addition, in distinction with Ref. \\cite{bghr}, we will not restrict ourselves to requiring $\\Omega_{\\tz_1}h^2<0.11$, but will instead allow all values of thermal neutralino abundance.  In Sec. \\ref{sec:method}, we provide an overview of our scanning methodology and parameter space values, and our calculation of the neutralino relic abundance. In Sec. \\ref{sec:wimp}, we present results of the neutralino relic density from a scan over  the SUGRA-19 model. We find that solutions with bino-like neutralinos give a large overabundance of CDM. Models with higgsino-like or wino-like neutralinos tend to give too little CDM, unless $m_{\\tz_1}\\sim 1$ TeV, with observable sparticles likely beyond LHC reach. The value $\\Omega_{\\tz_1}h^2\\sim 0.11$ lies at the minimum probability between these two extremes. Since our scan depends on the bias as to how we sample the GUT scale parameters, we present results for both a linear as well as a log scan; the latter scan accentuates small soft SUSY breaking parameter values. In Sec. \\ref{sec:sugra}, we present for comparison a similar analysis done for the mSUGRA model. In Sec. \\ref{sec:axino}, we work instead within the Peccei-Quinn augmented MSSM, or PQMSSM, which leads to mixed axion/axino CDM if the axino is assumed to be the lightest SUSY particle (LSP). The PQMSSM requires four additional parameters: the PQ breaking scale $f_a$ and initial mis-alignment angle $\\theta_i$, the axino mass $m_{\\ta}$, and the re-heat temperature of the universe $T_R$. In this case, the PQMSSM with 23 free parameters can more easily yield the measured relic abundance. We evaluate the favored range of the additional PQMSSM parameters. Our conclusions are presented in Sec. \\ref{sec:conclude}. ",
        "conclusions": "\\label{sec:conclude}  In this paper, we have reported on a calculation of the thermal abundance of relic neutralinos in the SUGRA-19 model with 19 free parameters, where all dimensional parameters are stipulated at the GUT scale, as suggested by gauge coupling unification, and SUSY GUTs\\cite{drw}. We find in this rather general framework that bino-like neutralinos populate the high $\\Omega_{\\tz_1}h^2\\sim 10^1-10^3$ region, while wino and higgsino-like neutralinos tend to populate the $\\Omega_{\\tz_1}h^2\\sim 10^{-3}-10^{-2}$ region. The measured CDM abundance sits at a deep minimum between these two favored values, especially if one requires rather light sparticles with $m_{\\tz_1}\\alt 500$ GeV, as suggested by electroweak fine-tuning arguments. In this region, a ``well-tempered neutralino'' is needed; however, these WTNs with the correct relic abundance are very difficult to generate using GUT scale SUSY breaking parameters. In this case, we conclude that it would be a near miracle if SUSY neutralinos were to constitute the measured CDM in the universe. This is the opposite conclusion to that suggested by the ``WIMP miracle'', where it is asserted that WIMPs have exactly the right properties to constitute thermal relics of the Big Bang. In the case of SUSY WIMPs, this does not appear to be so.  Our results were presented in the context of a linear scan over SUGRA-19 parameters, but remain qualitatively valid in the case of a log scan over parameter space. We also compare our results with those obtained in the minimal SUGRA or CMSSM model, where wino-like $\\tz_1$s don't appear. In the case of mSUGRA, the constrained GUT scale parameter choices actually make neutralino CDM somewhat more likely than in the SUGRA-19 case.  If the strong CP problem is solved by the PQ mechanism in the context of SUSY models, then an attractive alternative for dark matter appears: mixed axion/axino CDM. A scan over SUGRA-19 parameters, augmented by four extra PQMSSM parameters, shows that the mixed axion/axino DM abundance can be generated. While the measured CDM abundance doesn't lie at the peak of the predicted $a\\ta$ abundance, it also doesn't lie at a minimum. Therefore, in our viewpoint, the mixed axion/axino particles seem a more plausible candidate for CDM.  In the end, the issue will have to be resolved by experiment. A thorough search for relic axions with $f_a/N\\sim 10^9-10^{16}$ GeV is needed. Currently, only a tiny portion of QCD axion parameter space has been explored\\cite{admx}. A major consequence of our analysis then is to motivate our experimental colleagues to consider the possibility of new, more encompassing probes of axion dark matter. These should probe much more broadly in axion mass, and much more deeply in axion coupling."
    },
    "1005/1005.0026_arXiv.txt": {
        "abstract": "The discovery of the accelerated expansion of the universe using Type Ia supernovae (SNe Ia) has stimulated a tremendous amount of interest in the use of SNe Type Ia events as standard cosmological candles, and as a probe of the fundamental physics of dark energy. Recent observations of SNe Ia have indicated a significant population difference depending on the host galaxy. These observational findings are consistent with SNe Ia Ni-56 production in star-forming spiral galaxies  some 0.1 solar masses higher -- and therefore more luminous -- than in elliptical galaxies. We present recent full-star, 3D simulations of Type Ia supernovae which may help explain the nature of this systematic variation in SNe Ia luminosities, as well as the nature of the Ia explosion mechanism. These insights may in turn eventually shed light on the mystery of dark energy itself. ",
        "introduction": "The discovery of the accelerated expansion of the universe using Type Ia supernovae \\cite{riessetal98, perlmutteretal99} has stimulated a tremendous amount of interest in the use of SNe Type Ia events as standard cosmological candles, and as a probe of the fundamental physics of dark energy.  Supernovae come in different types, classified by their spectra. They are grouped into two main categories based on the presence or absence of Balmer absorption lines, indicating the presence or absence of atomic hydrogen. Type II contain hydrogen, whereas Type I lack hydrogen. These types are further subdivided by absorption line features and light curve shapes, see Table \\ref{table1}. In particular, Type Ia exhibit an absorption line of singly ionized silicon (Si$^+$) at 615 nm at peak light. Type Ib/c are further classified according to the absence of Si II lines, and the presence or absence of He lines, respectively.  Accordingly, Type Ia properties are consistent with white dwarf progenitors. The two leading models include white dwarfs in binary systems -- either with a non-degenerate companion, the so-called ``single-degenerate'' scenario, or with another white dwarf -- the so-called ``double-degenerate'' scenario. While both types of events are likely to occur in nature, and contribute to the diversity of explosion energies, there is still considerable debate over their relative frequencies, and even what the explosion mechanism is in each case. We outline some key theoretical and mounting observational evidence below based on nearby historic supernovae, which supports the identification of typical brightness ``Branch normal'' supernovae with single-degenerate events. \\begin{table} \\tbl{Classification of Supernovae} {\\begin{tabular}{ c c c } Mass $\\longrightarrow$ & & $\\longleftarrow$ Age \\\\ \\toprule SN Ia & SN II (IIn, IIL, IIP, IIb) & SN Ib/c \\\\ \\colrule No Hydrogen & Has Hydrogen & No Hydrogen \\\\ (Si$^+$ absorption) & & (no  Si$^+$) \\\\ White dwarf Progenitor & Core collapse of & Core collapse \\\\ & a massive star & (outer layers stripped by winds) \\\\ \\botrule \\end{tabular}} \\label{table1} \\end{table}    Tying the variable luminosities of Type Ia explosions to additional parameters is of key significance in any attempt to calibrate the Phillips relation. Analysis of light spectra demonstrates that Type Ia supernovae produce a relatively constant combined yield of stable nuclear statistical equilibrium  nuclei (NSE)  of  $^{58}$Ni and $^{54}$Fe and intermediate mass elements  (IME) Si-Ca of $1.05 \\pm .09\\ M_{\\odot}$ \\cite{mazzalietal07}. The approximately fixed burned mass over a wide range of SNe Ia luminosities suggests that SNe Ia share a single common explosion mechanism. However, recent observations of SNe Type Ia have indicated a significant population difference depending on the host galaxy. In particular, observations have found that SNe Ia in star-forming galaxies decline more slowly than those in ellipticals \\cite{howelletal07}.  These findings are consistent with   Ia SNe $^{56}$Ni production  in star-forming galaxies  some $\\sim 0.1\\ M_{\\odot}$ higher than that in ellipticals \\cite{pirobildsten08}.  Previous work has focused on the possible influence of the metallicity as a second parameter in explaining the luminosity variance of Type Ia events \\cite{timmesetal03, pirobildsten08, chamulaketal08, mengetal08}. These results suggest that while the metallicity effect may indeed contribute to the observed variance in Type Ia luminosities, alone it accounts for only a portion of the total variance \\cite {timmesetal03, pirobildsten08}. Consequently, we must look elsewhere to other physical effects -- including both the central  density and angular momentum profile of the white dwarf progenitors -- which may explain the majority of the observed Ia luminosity variance. These progenitor properties may in turn be directly influenced by the surrounding environment of Ia event, in particular the accretion rate from the companion star.  To put the single- and double-degenerate white dwarf progenitor models into their proper astrophysical context, we briefly consider the observational evidence. A number of recent papers have lended strong support for the general viability of the single-degenerate channel for the origin of Branch normal SNe Type Ia, in which a progenitor white dwarf accretes material from  a non-degenerate companion star \\cite{whelaniben73, nomoto82}. Analysis of the spectrum from the Type Ia SN 2006X in Virgo suggests that the blast wave moved through the circumstellar medium and collided with the ejecta from an otherwise undetected red giant companion star \\cite{patatetal07}. Furthermore,  a candidate G-type companion to SN1572 (Tycho)  has been identified \\cite{ruizlapuenteetal04}. SN1572 has independently been identified as a  Branch-normal Ia event  based on analysis of historical light curves and light echo spectra \\cite{ruizlapuente04, krauseetal08}.  However, despite this string of recent successes, the single degenerate model faces significant theoretical challenges. In particular, models suggest that a non-rotating white dwarf can burn stably only over a relatively narrow mass accretion rate range \\cite{nomoto82}. The addition of differential rotation broadens this range \\cite{yoonlanger04}. However, when one introduces rotation to the white dwarf structure, one must make a number of assumptions regarding accretion and internal shear in order to understand the evolution of the angular momentum distribution of the white dwarf progenitor. The accretion rate in turn sets the central density of the white dwarf at ignition.   As a consequence of these challenges faced by theoretical descriptions of the single degenerate model, and the limited direct observational evidence constraining progenitors, there is a significant degree of uncertainty in the characterization of the degenerate progenitor.  Here we focus on the role which the central density (or equivalently, the total mass) of a non-rotating white dwarf progenitor plays in determining the luminosity of Type Ia SNe. ",
        "conclusions": "The GCD model simulations demonstrate that successful self-consistent detonations of Type Ia supernovae in 3-D are possible without having to be initiated by hand.  The nuclear energy release during the deflagration phase is found to be sensitively dependent upon the central density of the white dwarf progenitor model.  With a greater pre-expansion afforded by the greater nuclear energy release in higher central density models, the GCD mechanism can lead to observed Branch normal luminosities, and typical levels of intermediate mass elements (Si -Ca)."
    },
    "1005/1005.2353_arXiv.txt": {
        "abstract": "$BVRIJHK$ photometry, {\\em Spitzer}-GLIMPSE photometry and $HK$ band spectroscopy were used to study the stellar content of IRAS 19343+2026, a (proto)star/cluster candidate, located close to the Galactic plane. The data suggest that IRAS 19343+2026 is a rich cluster associated with a massive protostar of 7.6 \\msun\\ with an age of $\\sim$ 10$^{5}$ yr. Three point sources in the vicinity of the far-infrared (FIR) peak are also found to be early B type stars. The remaining (predominantly low mass) members of the cluster are best represented by a 1 - 3 Myr pre-main-sequence (PMS) population. $HK$ band spectra of two bright and five faint point sources in the cluster confirm that the results obtained from the photometry are good representations of their young stellar object (YSO) nature. Thus, IRAS 19343+2026 is a young cluster with at least four early B-type stars classified as young (10$^{4}$ - 10$^{5}$ yr), that are surrounded by a somewhat older (1 - 3 Myr) population of low mass YSOs. Together, these results argue for a scenario in which low mass stars form prior to massive stars in a cluster forming environment. We compute the Initial Mass Function (IMF) for this cluster using the $K$-band luminosity function; the slope of the IMF is shallower than predicted by the Salpeter's mass function. The cluster mass, $M_{\\rm total}$, is estimated to be in the range $\\sim$ 307 \\msun\\ (from the data completeness limit) - 585 \\msun\\ (extrapolated down to the brown dwarf limit, assuming a certain IMF). ",
        "introduction": "One of the important issues in understanding massive star formation is the sequence of formation of low and high mass stars in a cluster forming environment \\citep{zb08}. This issue has relevance when explaining the nature of massive star formation in terms of either: (a) low to intermediate mass stars accumulating matter while evolving toward the main sequence, or (b) competitive accretion of massive cores in a star forming cloud (\\citealt*[]{mac02}; \\citealt*[]{mac03}; \\citealt*[]{zy07}).  Observations of embedded clusters generally display mass segregated configurations, with a few young high mass stars located at the cluster centers (\\citealt*[]{mcn86}, and references therein). Although an age estimation for the lower mass pre-main-sequence (PMS) population is possible using evolutionary tracks and colour-magnitude (CM) diagrams, the same is not true for massive stars, since they are usually only a few in number. Therefore, it is necessary to investigate clusters around high mass protostellar object candidates, where the youth of the massive stars is made clear by various other signposts, such as outflows, compact HII regions, excess FIR emission, and so on. \\citet{kum06} searched for clustering around a sample of high mass protostellar objects and found 54 embedded clusters among 215 candidates. The detection of clustering in about 25\\% of the sources (all targets located away from the Galactic mid-plane) suggests that at least one generation of low mass stars has already formed around these massive protostellar candidates. However, the Two Micron All Sky Survey (2MASS) data used by these authors did not effectively uncover the lower mass population; the 2MASS data were not deep enough to characterise the remaining cluster members. In contrast, a follow-up of the same data sets by \\citet{kum07} using the {\\em Spitzer}-GLIMPSE survey (\\citealt*[]{church09}) did not find any clustering for the sources in the Galactic mid-plane. This mid-infrared survey should have been more successful at finding clusters because of the better dust penetration.  In this paper, using high quality optical and infrared observations of a high mass protostellar object in the Galactic mid-plane, we attempt to examine such discrepancies, measure the stellar content, and understand the sequence of low and high mass star formation in the associated cluster. The high quality NIR photometric and spectroscopic data are combined with {\\em Spitzer} IRAC and MIPS observations to study the target. Study of the massive star content is particularly aided by radiative transfer modelling of the SED. The target, IRAS 19343+2026, is a high mass protostellar object candidate located at a kinematic distance (measured using NH$_3$ line velocities) of 4.2 kpc, with a FIR luminosity of 2.7$\\times$10$^4$ L$_{\\odot}$ \\citep{mol96}. This cluster is detected in the GLIMPSE survey (cluster 24 of \\citealt*{mcm05}).  Section 2 describes the detailed observations. The photometric, spectroscopic and SED modelling results are also presented in this section. The implications of the results are discussed in Section 3. We then summarize our conclusions in Section 4. ",
        "conclusions": "In the previous sections we characterised the stellar content associated with the massive protostellar object candidate, IRAS 19343+2026. The stellar content is composed of at least four young ($\\sim$ 10$^5$ yr) B type stars and a rich population of low mass stars at $\\sim$ 1- 3 Myr. The FIR peak is modelled as a 7.6 \\msun\\ massive young star. The cluster is surrounded by a bright infrared nebula seen from 3--24 $\\mu$m, indicative of dust and polycyclic aromatic hydrocarbons (PAH) emission. While the {\\em Spitzer} IRAC and MIPS bands display the infrared nebula that becomes increasingly brighter from 3.6 $\\mu$m to 24 $\\mu$m, the NIR bands predominantly display the dense stellar population embedded in this nebula. The following facts can explain why the cluster population is better detected in deep $K$ band images than in {\\em Spitzer} images. The contribution of stellar photospheres rapidly falls off longward of 2 $\\mu$m. Further, the background emission due to PAHs is strong in the bands longward of 4.5 $\\mu$m. The above two facts, together with the relatively large PSF of the {\\em Spitzer} data, reduces the sensitivity and contrast of the revealed low mass stars. However, the massive stars and the most embedded objects which appear bright at longer wavelengths are best revealed by the {\\em Spitzer} images. This can explain why deeper NIR $K$-band data such as the one presented here can unveil the lower mass population better than the GLIMPSE survey.  The fainter and relatively older low mass population revealed by the NIR images are uniformly arranged over the infrared nebula (see Fig. 1), coinciding well with the warm dust distribution seen in 24 $\\mu$m emission. As shown in Section 3.1, this population is best described to be 1 - 3 Myr old. In contrast, the four bright sources are found to be in the age range of 10$^4$ -- 10$^5$ yr (see Table 2). This suggests that low mass star formation occured in the cluster prior to the formation of massive stars. Although this idea has been argued before, from observations of other regions \\citep[e.g.][]{kum04}, we arrive at the same conclusion based on a more rigorous analysis using a CM diagram, $HK$-band spectra, and SED modeling. If all embedded clusters are born with a universal IMF, then sampling the IMF at any given time, uniformly over all mass ranges, is naturally expected to result in fewer massive stars and a greater number of low mass stars. Consequently, in the evolution of an embedded cluster (expected to follow an universal IMF), there exists an early phase when the total number of massive stars will be zero with a non-zero population of low mass stars. This simple statistical reasoning suggests that massive stars should appear after the low mass stars in a young cluster. Thus, it appears quite clear that the sequence for star formation is ``low mass stars first and massive stars next'', at least for clusters forming B type stars. The situation for clusters forming O type stars requires further investigation."
    },
    "1005/1005.3025_arXiv.txt": {
        "abstract": "{The well known radiogalaxy Cen\\,A has been recently detected as a source of very high energy (VHE) $\\gamma$-rays by the \\hess\\ experiment just before {\\it Fermi}/LAT detected it at high energies (HE). The detection, together with that of M87, established radiogalaxies as VHE $\\gamma$-ray emitters.} {The aim of this work is to present a lepto-hadronic model for the VHE emission from the relativistic jets in  FR I radiogalaxies.} {We consider that protons and electrons are accelerated in a compact region near the base of the jet, and they cool emitting multi wavelength radiation as propagating along the jet. The proton and electron distributions are obtained through steady-state transport equation taking into account acceleration, radiative and non-radiative cooling processes, as well as particle transport by convection.} {Considering the effects of photon absorption at different wavelengths, we calculate the radiation emitted by the primary protons and electrons, as well as the contribution of secondaries particles ($e^\\pm$, $\\pi$s and $\\mu$s ). The expected high-energy neutrino signal is also obtained and the possibility of detections with KM3NeT and IceCube is discussed.} {The spectral energy distribution obtained in our model with an appropriate set of parameters for an extended emission zone can account for much of the observed spectra for both AGNs.} ",
        "introduction": "The non-thermal high-energy emission from Active Galactic Nuclei (AGNs) has been widely studied in recent years at different wavelength ranges through both satellite-borne and ground-based detectors \\citep[e.g.][]{1995PASP..107..803U}. Several theoretical models have been proposed to explain the electromagnetic emission of these objects. It is commonly accepted that the high-energy radiation is emitted by particles accelerated in the relativistic jets launched from the inner parts of an accretion disk that surrounds the central black hole. In general, the high-energy spectral energy distribution (SED) of AGNs presents two characteristic bumps. The lower energy bump, located at optical to X-ray energies, is usually explained as synchrotron emission of electrons while the origin of the high-energy peak in the SED is still under debate {(see e.g. \\citep[e.g.][]{2007ASPC..373..169B} for a review)}. Leptonic models attribute this component to inverse-Compton up-scattering off synchrotron or external photons from the disk and/or radiation reprocessed in nearby clouds {\\citep[see e.g.][]{2001A+A...367..809K,2008A+A...478..111L,2005A+A...432..401G}}. In hadronic models, interactions of highly relativistic protons in the jet with ambient matter and photon fields, proton-induced cascades, or synchrotron radiation of protons, are responsible for the high-energy photons \\citep[see e.g.][]{2000CoPhC.124..290M,2001APh....15..121M,2004A+A...419...89R,2008AIPC.1085..644C,2009MNRAS.393.1041H,1996APh.....5..279R,2009AIPC.1123..242O}. There also exist models which are not based on the emission of accelerated particles in the relativistic jet and  assume the production of TeV $\\gamma$-rays in a pulsar-like cascade mechanism in the magnetosphere of the black hole \\citep[e.g.][]{2007ApJ...671...85N,2008A+A...479L...5R}.  The recently reported detection by HESS of the nearby radiogalaxy Cen\\,A  \\citep{2009ApJ...695L..40A} is of great relevance since it establishes radiogalaxies as VHE $\\gamma$-ray emitters. Cen\\,A is the second non-blazar AGN discovered at VHE, after the HEGRA detection of $\\gamma$-rays from M\\,87 \\citep{2003A+A...403L...1A} and the later confirmation by HESS \\citep{2006Sci...314.1424A}. A great variety of leptonic and hadronic models has been already applied to this kind of sources and {a full review is} beyond the scope of this work. During the first year of operation, {\\it Fermi}/LAT has detected HE emission from Centaurus\\,A \\citep{2009ApJ...707.1310A} and M\\,87 \\citep{2009ApJ...707...55A}, providing new constraints to the models.   In this work we present a lepto-hadronic model for the emission from FR I radiogalaxies. Section~\\ref{sec:sources} contains a brief description of observational facts on this type of sources. In Section~\\ref{sec:scena} we present the outline of our scenario, describing its most relevant characteristics. Section~\\ref{sec:model} is devoted to the description of the model. In Section~\\ref{sec:AplCenA} we present the application to Centaurus A, whereas in Section~\\ref{sec:AplM87} the results for M87 are given. Finally, in Section ~\\ref{sec:concl} we discuss the model implications and perspectives. ",
        "conclusions": "\\label{sec:concl}  We have presented here a lepto-hadronic model describing the particle propagation and interactions in the inclined jets of radiogalaxies. The scenario includes an acceleration zone placed near the jet base, where relativistic particles are injected. The location of this injection zone is fixed in the jet at a distance $z_{\\rm acc}$ from the black hole by requiring that the magnetic energy density is less prominent than the bulk kinetic density. This favors shock formation.  One very important aspect of the present model is that it does not consist of a one-zone treatment, since particles are allowed to convect along the jet, away from the injection zone. To obtain the particle distributions, we have used a stationary one-dimensional transport equation that accounts for energy losses, escape, and convection of particles. The resulting distribution of primary electrons is important mainly in the injection zone, since they undergo a rapid synchrotron cooling. Instead, protons are convected away along the jet, loosing energy gradually (mainly by adiabatic cooling, synchrotron radiation, and $pp$ interactions). Also they undergo significant $p\\gamma$ collisions with the electron-synchrotron photons at the injection zone. The secondary pions and muons produced in the hadronic interactions are also considered, and their distributions are found using the mentioned type of transport equation. This allows to account for the synchrotron cooling suffered before the decay of this secondary, transient particles.  We have found that the radiative output obtained can account for much of the observational data for the two radiogalaxies M87 and Cen\\,A. As seen in Figs.~\\ref{SED_cena} and ~\\ref{sed-M87}, the model allows for the possibility of explaining the VHE emission by means of $pp$ interactions. The hard X-ray data corresponds to electron synchrotron radiation whereas the soft-$\\gamma$ emission is well reproduced by proton synchrotron and IC emission. {Previous hadronic models (e.g. the synchrotron proton blazar of \\cite{2003APh....18..593M}) could also fit the Fermi energy range for M87 and the hard spectrum measured by HESS (consistent with VERITAS and MAGIC measurements), according to what it is mentioned in \\cite{2009ApJ...707...55A}}. A previous one-zone lepto-hadronic model for Cen\\,A was proposed by \\cite{2009AIPC.1123..242O} just before HESS and Fermi/LAT data was available.  In the present work, we have included the effects of absorption due to photoionization interactions in the surrounding {medium}. This causes a drastic modulation in the electron synchrotron spectrum, especially for Cen\\,A, which is responsible for the whole emission in the broadband range $\\sim 10^{-5}-10^{7}$eV. In the case of M87, this type of absorption has a minor impact, since the column density is much less ($N_H=2 \\times 10^{20}{\\rm cm}^{-2}$, \\cite{1996ApJ...458L...5L}). The internal absorption of gamma-rays is also taken into account, being important only within the injection zone, where the synchrotron emission of primary electrons provide an important absorbing target.  It is important to note that the same model can be used to describe different types of jet. In terms of the broadband photon emission, M87 data shows a similar luminosity level for high and the low energy ranges. Since the VHE emission is basically determined by $pp$ interactions, we need a large power to be injected in relativistic protons as compared to that injected in electrons. This is because electrons cool completely, mainly by synchrotron emission, and protons undergo an important adiabatic cooling. Hence, in the case of M87, we need a high value for the proton-to-electron power ratio: $a=40$. A different situation arises for Cen A, where the highest luminosities correspond to energies $E_{\\gamma}<10^7{\\rm eV}$, and the HE and VHE luminosities are comparatively lower. In this case, we need {more power to be injected in electrons than in protons}  ($a=2.5\\times 10^{-2}$) to allow the broadband electron-synchrotron spectrum to approximately account for all the data for $E_{\\gamma}<10^7{\\rm eV}$. Then, for Cen\\,A a lower luminosity in protons is enough to reach the level of the observed HE and VHE radiation through proton-synchrotron and $pp$ interactions, respectively. A further difference between Cen\\,A and M87 is in the slope of the injected particle distributions: for Cen\\,A, we have a quite flat injection ($s=1.8$), while for M87 we find a steeper injection ($s=2.4$). Also the acceleration efficiency is greater for Cen A than for M87. This also has an impact in the neutrino spectrum produced, which for $E_\\nu> 1 $TeV appears more difficult to be detected for M87 than for Cen\\,A.  Finally, we remark that a possible improvement of the present treatment, which is left for future work, should be the inclusion of a time-dependent injection to account for flaring states."
    },
    "1005/1005.3814_arXiv.txt": {
        "abstract": "In the coming decade, low-frequency radio arrays will begin to probe the epoch of reionization via the redshifted 21-cm hydrogen line. Successful interpretation of these observations will require effective statistical techniques for analyzing the data. Due to the difficulty of these measurements, it is important to develop techniques beyond the standard power spectrum analysis in order to offer independent confirmation of the reionization history, probe different aspects of the topology of reionization, and have different systematic errors. In order to assess the promise of probability distribution functions (PDFs) as statistical analysis tools in 21-cm cosmology, we first measure the 21-cm brightness temperature (one-point) PDFs in six different reionization simulations. We then parametrize their most distinct features by fitting them to a simple model. Using the same simulations, we also present the first measurements of difference PDFs in simulations of reionization. We find that while these statistics probe the properties of the ionizing sources, they are relatively independent of small-scale, sub-grid astrophysics. We discuss the additional information that the difference PDF can provide on top of the power spectrum and the one-point PDF. ",
        "introduction": "The first stars and quasars reionized the intergalactic medium (IGM) during the epoch of reionization (EOR) \\citep{Review}. From the Ly$\\alpha$ absorption in the spectra of distant quasars, we know that the EOR ended at $z\\sim 6$ \\citep[see, e.g., ][]{dj05}, and from the cosmic microwave background (CMB) polarization maps we can infer that it started no later than $z\\sim 15$ \\citep{WMAP}. Understanding how the reionization process developed over time offers a way to answer some of the critical questions of modern cosmology concerning properties of the first sources of light in the universe. Among the most promising observational probes of the EOR is the 21-cm spectral line, from hyperfine splitting in the ground state of hydrogen, with an energy of $5.9\\times 10^{-6}$eV that corresponds to the rest-frame frequency of 1420 MHz. The redshifted 21-cm emission from neutral regions of the IGM during reionization is estimated to be a $1\\%$ correction to the energy density of the CMB. It is expected to display angular structure and frequency structure, due to the inhomogeneities in the gas density, ionized fraction, $x_i$\\footnote{The ionized fraction $x_i$ is the ratio of the number of protons to the total number of hydrogen atoms.}, and spin temperature \\citep{md97} of the emitting gas.  Statistical detection of the large-scale brightness fluctuations is within the scope of a number of experiments that are presently being built, such as the Murchison Widefield Array (MWA) and the Low Frequency Array (LOFAR)\\citep[for reviews see, e.g.,][]{fob06, bl07}. In this context, it is important to develop appropriate statistical tools to be employed in analyzing the incoming data. Such development is facilitated by the fact that the N-body and radiative-transfer simulations of reionization have begun to reach the large scales of order 100 comoving Mpc \\citep{mq07,iliev,santos} needed to capture the evolution of the IGM during the EOR. These simulated data cubes can be used to test various statistical tools proposed for extracting information about the properties of the IGM during reionization.  So far, studies of the statistics of the 21-cm fluctuations have mainly focused on the power spectrum of the brightness temperature, $T_b$ \\citep{miguel, CfA}. While this statistic is fully representative at the onset of the EOR, where the Gaussian primordial density fluctuations drive the 21-cm fluctuations, it ceases to be so at later times. Namely, as the reionization process advances, the mapping between the hydrogen density and $T_b$ becomes highly non-linear (as evidenced for instance by the bounded domain, $x_i \\in [0,1]$), which results in non-Gaussianity of the probability distribution function (PDF) of $T_b$.  For this reason, various authors have started exploring alternative and complementary statistics \\citep[e.g., ][]{fzh04, Ali}, in particular the PDFs and difference PDFs of the 21-cm brightness temperature \\citep{dif08, ich09}. We test these two statistics on six different simulated data cubes from \\citet{mq07}. These data cubes are results of different astrophysical inputs that produce various reionization histories, all of which are allowed by the current observational constraints. We measure the one-point PDFs and difference PDFs and analyze their properties.  The plan of the paper is as follows: in Section~\\ref{sims} we briefly describe the simulation runs used in this paper. In Section~\\ref{onept} we then present the measured one-point PDFs along with the best fits of the model proposed by \\citet{ich09}, and discuss the main parameters driving the PDF shape. We next present in Section~\\ref{diff} the first measurements of difference PDFs for the same set of simulations and analyze their properties. We conclude in Section~\\ref{conclude}. ",
        "conclusions": "Upcoming low-frequency radio observations will use the MWA, LOFAR, and similar instruments to survey the sky for redshifted 21-cm emission. It is therefore important to develop statistical tools that can be used to extract information about the history of reionization from these observations. In this paper, we examine PDFs and difference PDFs, using six different simulations from \\citet{mq07}. We show that the PDFs are highly non-Gaussian in the midst of the reionization process, and are thus a complementary statistic to the commonly discussed power spectrum. As a way to analyze the PDFs, we examine the evolution of the parameters of the best-fit GED model with the global ionized fraction $\\bar{x}_i$. In particular, the $\\delta_D$ function portion of the probability ($P_D$) measures a low-resolution, smoothed-out version of the reionization history (i.e., $\\bar{x}_i$ as a function of redshift).  We also present the first numerical measurements of difference PDFs; specifically, we measure difference PDFs for the same set of simulation runs from \\citet{mq07}. We argue that the larger data set and the nature of this statistic can be significant advantages in the presence of bright foregrounds. The difference PDF can be physically understood as arising from three contributions: pixel pairs in which both, one, or neither of the pixels is ionized. As an illustration of the information that can be deduced from the difference PDFs, we consider the typical correlation length, which corresponds to the average size of an ionized bubble during reionization, or the typical density fluctuation correlation length, at the onset of the EOR. The difference PDF also has a delta-function portion ($\\Delta P_D$), which measures a low-resolution, smoothed-out version of the ionization correlation function at each redshift.  We find that increasing small-scale clumping, and including photoheating feedback or minihaloes has only a small effect on the one-point and difference PDFs (considered at a given $\\bar{x}_i$), at least within the range of assumptions covered by the simulations that we considered. On the other hand, the PDFs are highly sensitive to the properties of the ionizing sources, so that measuring them can help distinguish between reionization driven by large versus small haloes and help us unveil information about the first sources of light in the universe. These conclusions parallel those of \\citet{mq07}, highlighting the fundamental fact that the spatial structure of reionization is driven by large-scale modes and depends mainly on the overall bias of the ionizing sources \\citep{bl04,fzh04}.  We plan to more precisely quantify the properties of difference PDF and establish the relation between their features and the properties of the IGM during the EOR. It would also be interesting to explore the PDFs and difference PDFs in alternative reionization scenarios, such as those dominated by X-ray sources or a decaying dark matter particle. Contrasting different scenarios may lead to a fuller understanding of the information content of the specific PDF shapes that we measured."
    },
    "1005/1005.4928_arXiv.txt": {
        "abstract": "Although rapid reorientation of a black hole spin axis (lasting less than a few Myr) has been suggested as a mechanism for the formation of wings in X-shaped radio galaxies (XRGs), to date no convincing case of reorientation has been found in any XRG.  Alternative wing formation models such as the hydrodynamic backflow models are supported by observed trends indicating that XRGs form preferentially with jets aligned along the major axis of the surrounding medium and wings along the minor axis. In this Letter, we present a deep {\\em Chandra} observation of 4C~+00.58, an oddball XRG with its jet oriented along the {\\em minor} axis.  By using the X-ray data in tandem with available radio and optical data, we estimate relevant timescales with which to evaluate wing formation models. The hydrodynamic models have difficulty explaining the long wings, whereas the presence of X-ray cavities (suggesting jet activity along a prior axis) and a potential stellar shell (indicating a recent merger) favor a merger-induced reorientation model. ",
        "introduction": "X-shaped radio galaxies (XRGs) are double-lobed radio galaxies \\citep{leahy84} which also possess a pair of long, faint, centro-symmetric ``wings''. They have gained notoriety as a possible signature of a rapid (within a few Myr) reorientation of the supermassive black hole (SMBH) spin axis, conceivably induced by galaxy mergers in which either accretion torque or a SMBH merger causes a spin-flip \\citep{rottmann01,merritt02}.  In this scenario, wings are ``fossils'' tracing the prior jet axis which radiatively decay.  However, no convincing case for reorientation has been made in any individual XRG, whereas several lines of evidence support a hydrodynamic origin for the wings.  For instance, in most XRGs, the wings are coaligned with the minor axis of the host galaxy and the jet with the major axis \\citep{capetti02,saripalli09}. A similar major-axis---radio alignment trend exists in the X-rays \\citep{hodges10}.  These results have been interpreted to favor models in which the wings are produced by radio-lobe---gas interaction \\citep{leahy84,worrall95,capetti02}.  Additionally, no clear signs of mergers have been found in a spectroscopic study of XRG hosts, whereas they may be overpressured \\citep{landt10}.  In this Letter, we identify 4C~+00.58 \\citep[Fig.~1, classified as a candidate XRG by][]{cheung07} as one of the best candidates for a merger-induced reorientation based on quantities derived from a deep {\\it Chandra X-ray Observatory} observation and publicly available data.  Unlike other XRGs, the 4C~+00.58 jet is coaligned with the {\\it minor} axis of its host. Even if only a fraction of XRGs are produced by reorientation, their frequency may be important for estimating {\\it Laser Interferometer Space Antenna} detection rates.  We use a Galactic absorption of $N_H = 7.14\\times10^{20}$~cm$^{-2}$ \\citep{kalberla05}, as well as the {\\em Wilkinson Microwave Anisotropy Probe} cosmology \\citep[$H_0 = 71$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\Lambda} = 0.73$, and $\\Omega_m = 0.27$;][]{spergel07}.  At a redshift of $z = 0.059$, $1^{\\prime\\prime} = 1.13$~kpc. ",
        "conclusions": ""
    },
    "1005/1005.2435_arXiv.txt": {
        "abstract": "We present large scale 3D particle-in-cell (PIC) simulations to examine particle energization in magnetic reconnection of relativistic electron-positron (pair) plasmas. The initial configuration is set up as a relativistic Harris equilibrium without a guide field. These simulations are large enough to accommodate a sufficient number of tearing and kink modes. Contrary to the non-relativistic limit, the linear tearing instability is faster than the linear kink instability, at least in our specific parameters. We find that the magnetic energy dissipation is first facilitated by the tearing instability and followed by the secondary kink instability. Particles are mostly energized inside the magnetic islands during the tearing stage due to the spatially varying electric fields produced by the outflows from reconnection. Secondary kink instability leads to additional particle acceleration. Accelerated particles are, however,  observed to be thermalized quickly. The large amplitude of the vertical magnetic field resulting from the tearing modes by the secondary kink modes further help thermalizing the non-thermal particles generated from the secondary kink instability. Implications of these results for astrophysics are briefly discussed. ",
        "introduction": "Magnetic reconnection in collisionless relativistic electron-position (pair) plasma plays an essential role in a number of important astrophysical problems, such as relativistic jets from supermassive black holes, relativistic winds from pulsars and possibly outflows from Gamma-ray Bursts. In the Crab pulsar wind, relativistic reconnection in pair plasmas has been considered as a primary mechanism to convert magnetic energy to particle energy. \\cite{ks03,zh05} Although magnetic reconnection has been extensively studied, the physics of fast reconnection and particle energization in 3D relativistic reconnection is still a mystery. Magnetic reconnection of pair plasmas provides a unique opportunity to examine critically the fundamental reconnection physics due to the absence of Hall term, \\cite{bb05} which is generally thought to be the main mechanism for fast reconnection in electron-ion plasmas. Understanding of relativistic magnetic reconnection of pair plasmas will also provide useful physical insights for both new laboratory experiments of pair plasmas \\cite{cwb09} and observations of the high energy Universe.  Previous reconnection studies of relativistic pair plasmas have focused on the fast reconnection and particle acceleration mostly in 2D \\cite{zh01, zh05, bb07, zh07} or reconnection onset in 3D with relatively small system size. \\cite{jtls04, zh05b, zh08}  It has been shown that instabilities in a cross-field plane such as the Kelvin-Helmholtz instability (KHI) \\cite{yda96,ssf01} and the drift kink instability (DKI) \\cite{dw98,zh05} are of critical importance to understanding magnetic energy dissipation. \\citep{zh05} The past studies used small-scale particle-in-cell (PIC) simulations, (e.g., the length in the outflow direction is $\\sim25$ inertial lengths), however, the characteristic system size in most astrophysical/laboratory systems is vastly larger. Thus, it is important to study these processes in a truly 3D configuration. Fast reconnection in the relativistic limit in 2D has been reported to be mostly facilitated by pair pressure tensors. \\cite{hz07,bb07} \\citet{ydk08} performed the largest 3D PIC simulations to date of non-relativistic pair plasma reconnection and found that elongated current sheet is unstable to secondary kink instability and leads to the bending of the current sheet. For particle acceleration, past 2D simulations of reconnection mostly started from an initially externally imposed reconnection electrical field and showed that particles are primarily accelerated near the X lines by electrical fields. \\cite{hmts01,pp06,zh01,zh07} However, electrical fields within a small fraction of the total volume at X lines are difficult to account for the large number of accelerated electrons of the total volume. \\cite{cbp08} Recently \\citet{cbp08} reported some surprising satellite observation evidences that energetic electron flux peaks at sites of compressed density within islands during electon-ion reconnection in the Earth's magnetosphere. \\citet{hz07} reported that for relativistic pair plasmas in 2D without a guide field, most of particle acceleration happens inside magnetic islands possibly like a betatron (but no details are provided) (also see \\citet{dms04}). Thus, the well accepted particle acceleration mechanism needs to be revisited, especially in 3D. \\citet{dscs06} suggests that Fermi acceleration from contracting magnetic islands is the major mechanism for acceleration in 3D non-relativistic electron-ion plasma magnetic reconnection and electrons are mostly accelerated within the islands. In this paper, we will present the first numerical study of 3D magnetic reconnection of relativistic pair plasma with an emphasis on particle energization.  The paper is organized as follows. Section~\\ref{setup} describes the computational model and problem setup.  In Sec.~\\ref{results}, overall simulation evolution is briefly described followed by particle energization and its mechanisms, particle thermalization and the role of the secondary kink instability on particle acceleration in magnetic reconnection. Conclusions and their implications for astrophysical situations are given in Sec.~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  In this paper we have presented large scale 3D particle-in-cell (PIC) simulations to examine particle energization in magnetic reconnection of relativistic electron-positron (pair) plasmas. These simulations are large enough to accommodate a sufficient number of tearing and kink modes. The interplay of these two main instabilities regulates the overall evolution: from the initial linear tearing and kink growth with the appearance of small magnetic islands, to the merging of small magnetic islands into a big one, establishing the major reconnection sites, to the elongated current sheet becoming unstable to the secondary kink instability, forming plasmoids. The particle energization could be divided into two main stages: (1) the tearing stage ($t\\Omega_c\\lesssim954)$, and (2) the secondary kink stage ($t\\Omega_c\\gtrsim954)$. We find that particles are mostly energized inside the magnetic islands during the tearing stage due to the spatially varying electric fields produced by the outflows from reconnection. Secondary kink instability also leads to some particle acceleration. Accelerated particles are, however,  observed to be thermalized. Since the physical conditions for both the energization process and the thermalization process discussed in this study are quite typical in relativistic pair reconnection, we suggest that these processes are germane in understanding the magnetic energy dissipation in pair plasmas.  In a relativistic current sheet, previous literature reported that the linear kink instability is faster than the linear tearing instability unless the reconnection is driven.\\cite{zh05,zh05b,zh07,zh08} In this work, contrary to the above well-known results, the linear tearing instability is faster than the linear kink instability, at least in our specific parameters. This result raises another hint to the striped wind problem.\\cite{ks03}  It is worth of emphasizing that all the calculations reported in this paper are done with an initial anti-parallel magnetic field without a guide field for pair plasmas. The introduction of a uniform guide field would significantly change the results as reported in \\citet{hz07}. Compared to a relativistic reconnection without a guide field, the reconnection rate is slower with a guide field due to a lower compressibility effect of the guide magnetic field. More importantly with a guide field, the particle acceleration is mostly done at the reconnection site, whereas it is inside the magnetic island without a guide field.\\cite{hz07} The magnetic field configuration in the real astrophysical situation such as pulsar wind would be even more complicated. We need to be cautious when trying to apply the particle energization processes reported in this paper to the real astrophysical situations. Additionally the energization process in pair plasmas may find only limited applications in the  solar or magnetospheric environments  where electron-ion reconnection should be dominant."
    },
    "1005/1005.0590.txt": {
        "abstract": "One of the greatest problems of standard cosmology is the Big Bang singularity. Previously it has been shown that non-local ghostfree higher-derivative modifications of Einstein gravity in the ultra-violet regime can admit non-singular bouncing solutions. In this paper we study in more details the dynamical properties of the equations of motion for these  theories of gravity in presence of positive and negative cosmological constants and radiation. We find stable inflationary attractor solutions in the presence of a positive cosmological constant which renders inflation {\\it geodesically complete}, while in the presence of a negative cosmological constant a cyclic universe emerges. We also provide an algorithm for tracking the super-Hubble perturbations during the bounce and show that the bouncing solutions are  free from any perturbative instability. ",
        "introduction": "Primordial {\\it inflation} is one of the most well motivated paradigms of the early universe, for a recent review, see~\\cite{MR}. A large e-foldings of inflation can explain the large scale structures of the universe and also the origin of seed perturbations for the cosmic microwave background radiation. In spite of the great successes of inflation, it does not address one of the most important aspects of the Big Bang cosmology. For any equation of state obeying the strong energy condition $p>-\\rho/3$, regardless of the geometry (flat, open, closed) of the universe, the scale factor of the universe in a Friedmann Robertson Walker (FRW) metric vanishes at $t=0$, and the matter density diverges. In fact all the curvature invariants, such as $R,~\\Box R,...$, become singular. This is the reason why it is called the {\\it Big Bang singularity problem}.  Although inflation requires $p<-\\rho/3$, it does not alleviate the Big Bang singularity problem, rather it pushes the singularity backwards in time. Many authors have pondered on this debatable issue, whether inflation is past eternal or not~\\cite{Borde,Guth,Linde}, and the conclusion is that it is not, at least in the context of Einstein gravity as long as the average expansion rate in the past is greater than zero, i.e. $H_{av}>0$~\\cite{Guth}. The fluctuations grow as the universe approaches the singularity, and the standard {\\it singularity theorems} due to Hawking and Penrose hold, which inevitably leads to a collapse in FRW geometry as long as the energy density is positive~\\cite{Hawking} (see also~\\cite{Guth}).  There have been many attempts to resolve this cosmological singularity problem classically~\\cite{Hawking}, and in the context of string theory~\\cite{Many}, for a review see~\\cite{Craps}, but none has successfully resolved the space-like singularity in the context of a flat universe. All the stringy toy models which have been constructed so far seem to suffer from classical and/or quantum instabilities and inconsistencies, %in the presence of time dependent background such as the presence of negative tension branes, closed time like curves, negative-norm ghost states, etc., see~\\cite{Cornalba,Horowitz,Seiberg} for details.  One alternative view would be to study fully non-perturbative and ghostfree higher derivative actions of gravity inspired from string  theory~\\cite{BMS}: An infinite set of higher derivative corrections (series expansion in $\\al'$, the string tension) mostly in the form of an exponential are often known to appear in stringy models, such as tachyonic actions in string field theory~\\cite{sft1}-\\cite{pressure} (for a review see~\\cite{sft_review}),  Bulk fields localized on a codimension-2 branes \\cite{claudia} and various toy models of string theory such as $p$-adic strings~\\cite{padic_st} (see also~\\cite{zwiebach}), zeta strings~\\cite{zeta}, and strings quantized on a random lattice \\cite{random,marc} (see also \\cite{ghoshal}). These corrections appear already at classical level (i.e., at the tree level), as well as  loop level or even non-perturbatively (e.g., as expected in the 1/N expansion in some Yang-Mills theory~\\cite{Hooft}). Incidentally, recently it was shown~\\cite{bck} how one can consistently perform quantum loop calculations in these non-local models  and  recover, in the context of finite temperature field theory, some of the novel properties of string theory such as thermal duality and Hagedorn transition. Thus it is interesting to explore whether similar theories of gravity, albeit only at the level of toy models, can provide us insights into the resolution of cosmological singularities.  It is worth while to point out that like most higher derivative actions, these theories have improved Ultra-violet (UV) behavior and  can render gravity asymptotically free, but importantly without introducing ghost states~\\cite{BMS}, for some earlier attempts, see~\\cite{Stelle,Van,Zwiebach}, which were not very successful.  Note that asymptotic freedom  is required (in addition to renormalizability) in ordinary field theory for phenomenological reasons (scaling in processes such as deep inelastic scattering), as well as to avoid certain problems in the nonperturbative definition of the theory. This condition usually takes the form of ``asymptotic safety\" \\cite{Weinberg}, a more general requirement for having a well-defined UV limit. The theories we consider are therefore natural candidates for such GR extensions. Building on the work done in~\\cite{BMS}, in this paper we continue to study cosmological  background solutions and their perturbations arising in the  higher derivative nonlocal gravity models, and in particular whether  they can resolve the Big Crunch/Bang singularity problem of GR robustly. Previously, non-local~\\footnote{Some of these models are very similar to the models we consider in our paper and contain and infinite series of higher derivative terms, but in the literature highly non-local operators, such as $f(1/\\Box)$, have also been considered.}   higher derivative corrections to gravity has been studied in the context of black hole singularities~\\cite{warren}, understanding ultra-violet behavior~\\cite{uv-nonlocal},  cosmological applications (for instance, involving  dark energy~\\cite{Woodard},  inflation~\\cite{inflation}, and  generating density perturbations~\\cite{pert}), and modifications to the Newtonian potentials in the metric~\\cite{nonlocal-gravity}. For discussions on these topics in the more general context of higher derivative theories of gravity, see~\\cite{gravity}~\\footnote{In parallel, string-inspired non-local scalar field theories have also been studied quite extensively in recent times, including  several interesting applications~\\cite{scalar}.}.  In~\\cite{BMS}  non-local asymptotically safe theories of gravity were already constructed and it was shown that they can admit non-singular bouncing solutions. However, it was not clear how generic these solutions were, whether these solutions were stable under small perturbations, and whether these theories contain other singular solutions or not. The primary focus of this paper is to shed some light on these issues, and also see whether they can provide phenomenologically viable cosmological models. For instance, can one have a geodesically complete inflationary paradigm once the non-local modifications to gravity are included? We know that for the success of the inflationary paradigm,  the super-Hubble fluctuations must ``freeze'', is that true in the non-local models or do the fluctuations decay/ grow  during inflation? In cyclic and bouncing alternatives~\\cite{narlikar}-\\cite{prebigbang}, the issue of transmitting perturbations is of paramount importance. Can we understand the evolution of the fluctuations in the concrete nonsingular bouncing set-up that our model provides?  In this paper specifically we will seek a two parameter family of solutions, labeled by $\\lambda$ and $\\mu$, which solves the higher derivative theory of gravity. Interestingly, this class of solutions demands the presence of a cosmological constant, which is no longer free but is determined in terms of the two parameters, and  radiation. One particular solution for the scale factor turns out to be a {\\it hyperbolic cosine } bounce with a positive cosmological constant, where $\\mu=-6\\lambda^2$~\\cite{BMS}. In principle $\\lambda$ can be both positive and negative in order to provide a non-singular bouncing cosmology, but for $\\lambda < 0$, it is not possible to recover the GR limit at late times. The most interesting case is when $\\lambda >0$ where not only we recover the right infrared behavior of gravity, but also two interesting classes of solutions with positive and negative cosmological constants, i.e. $\\pm \\Lambda$.   The case of $\\Lambda >0$ is very interesting from the point of view of inflation, as it paves the way for a {\\it geodesically} complete paradigm in past and future. There exists a deSitter solution as an attractor for large enough times before and after the bounce. This is a good news for both low scale~\\cite{AEGJM} and an intermediate scale models of inflation~\\cite{EMS} which are driven by the vacuum energy of the inflaton field which can excite the observed relativistic and non-relativistic degrees of freedom after the end of inflation. On the other hand the $\\mu>0,\\Lambda <0$ case is equally interesting from the point of view of providing an early phase of  cyclic cosmology. In this case there does not exist any deSitter attractor but there are also no singular solutions. Although, a mere cyclic phase is a rather boring scenario as it fails to reproduce the universe we live in,  recently it has been proposed that once interactions between radiation and non-relativistic matter are taken into account then we can obtain a cyclic inflationary paradigm where the  cycles become asymmetric and the universe expands bit by bit  to give rise to an exponential inflation over many many cycles~\\cite{BM}, see also~\\cite{AB}. Both $\\Lambda>0, \\Lambda<0$ cases will be studied analytically and numerically in our paper.  In general, taking into account a preceding contracting phase removes the inflationary infrared divergences of quantum fluctuations \\cite{Koivisto:2010pj}. Here we study the evolution of the super-Hubble perturbations during the bounce and ensure that they do not feature any classical and/or quantum instabilities. Interestingly, we find that on large scales all the modes are decaying except one which remains frozen. This frozen mode could be useful if we wish to evolve the perturbations either prior or after the bounce in order to match the observed CMB perturbations. However, we will not delve into any details about these issues in this paper.  We will begin our discussion with a brief recap of nonperturbative  higher derivative action and its equation of motion in section 2. We will discuss the background solutions in section 3, where we also provide an analytical example of a specific class of non-local models. In section 4, we will study various cases of $\\lambda >0, \\lambda <0$  and their dynamical attractor solutions for $\\Lambda >0,\\Lambda <0$. In section 5, we track the perturbations during the bounce. Some series expansions are confined to the appendices: \\ref{app1} presents a general series solution for the two-parameter ansatz, and \\ref{app2}, \\ref{app4} \\ref{app3} consider some special cases. We conclude in section 6.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "In this work we have shown how ghost and asymptotically free non-perturbative effective actions of gravity motivated from string theory are able to realize  non-singular bouncing cosmologies. We studied a class of solutions with two parameters, which solves the non-local and non-linear equations of motion for gravity in presence of a cosmological constant and radiation energy density. In particular, a specific class of non-singular bouncing solutions is given by a hyperbolic cosine scale factor of the universe.  We studied cosmological scenarios with both positive and negative cosmological constants. In the former case we have found that an inflationary (deflationary) solution arises as a late time attractor  after (before) the bounce, the non-local contributions to  gravity vanishes, and the space-time asymptotes towards the GR deSitter limit. This result is very relevant for inflationary cosmology as now time-like and null-like trajectories can be made {\\it geodesically complete} within the inflationary framework. Without the non-singular bounce the trajectories in the past light cone would always be finite, which would render inflation geodesically past incomplete. Furthermore, the Bunch-Davies vacuum would be pathological in the infrared. %We also argued that in order to get an inflationary attractor the cosmological constant is bounded from below, i.e. $\\Lambda \\geq (M_{\\ast}^2M_p^2)/8$, where $M_*$ is the scale at which the higher derivative operators in the action become important.  When the cosmological constant becomes negative an interesting stable solution arises which leads to a cyclic universe in presence of radiation. Although just with radiation a cyclic universe does not lead to any phenomenologically viable  cosmology,   exciting cosmology can arise if we add a non-relativistic matter component and interactions between radiation and matter. As shown in \\cite{BM}, such an interaction would lead to a cyclic inflationary scenario with distinct signatures for the cosmic microwave background perturbations. Also, our non-local model provides a concrete setting to study how perturbations are transmitted through a nonsingular bounce which is important for cyclic and bouncing models in general. Finally, the non-local actions with $\\la>0,\\La<0$ and $\\mu<3\\la^2$, constitutes an especially nice class of actions because they don't seem to contain any singular solutions (at least within the confines of our ansatz), which is not the case for actions with $\\la,\\La>0$. An important task for the future is to determine whether these non-local actions admit any singular solutions at all or not.  In order to make sure that both inflationary and cyclic attractors are stable, and there exists no instabilities during the bounce, we studied the long wavelength, super-Hubble perturbations. We found that all the modes of the long wavelength perturbations are decaying except the one which is frozen. The presence of a constant mode is a good news for seeding density perturbations for the cosmic microwave background radiation during, before or after the bounce. The perturbation analysis suggests that there are no pathologies at least on large scales which would affect the non-singular bounce  solution. Studying the sub-Hubble perturbations is more difficult in higher derivative theories of gravity. However, this is necessary to consider the ghost problem in the bouncing backgrounds, whereas we have yet only shown that the theories are not plagued by this problem in Minkowski backgrounds. We hope to revisit this issue in a separate publication.   To summarize, we have found  non-singular bouncing solutions for the scale factor of the universe in a special class of nonlocal gravitational models which are ghost and asymptotically free in the ultra violet, but nevertheless recover the GR limit at large distance and time scales. We also made preliminary progress to study perturbations around the bouncing background and in particular demonstrated that the bouncing solutions are late time attractors and perturbatively stable.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1005/1005.0740_arXiv.txt": {
        "abstract": "% \\cyg\\ is a microquasar consisting of an accreting compact object orbiting around a Wolf--Rayet star. It has been detected at radio frequencies and up to high-energy $\\gamma$ rays (above 100 MeV). However, many models also predict a very high energy (VHE) emission (above hundreds of GeV) when the source displays relativistic persistent jets or transient ejections. Therefore, detecting such emission would improve the understanding of the jet physics. The imaging atmospheric Cherenkov telescope MAGIC observed \\cyg\\ for about 70 hr between 2006 March and 2009 August in different X-ray/radio spectral states and also during a period of enhanced $\\gamma$-ray emission. MAGIC found no evidence for a VHE signal from the direction of the microquasar. An upper limit to the integral flux for energies higher than 250 GeV has been set to 2.2~$\\times$~$10^{-12}$~$\\mathrm{photons}$~$\\mathrm{cm^{-2}}$~$\\mathrm{s^{-1}}$ (95$\\%$ confidence level). This is the best limit so far to the VHE emission from this source. The non-detection of a VHE signal during the period of activity in the high-energy band sheds light on the location of the possible VHE radiation favoring the emission from the innermost region of the jets, where absorption is significant. The current and future generations of Cherenkov telescopes may detect a signal under precise spectral conditions. ",
        "introduction": "\\cyg\\ is a bright and persistent X-ray binary, discovered in 1966~\\citep{Giacconi67}, but the high-energy processes occurring in the source are still poorly understood. It lies close to the Galactic plane at a distance between 3.4 and 9.8~kpc, probably at 7~kpc, depending on different distance estimates to the Cygnus OB2 association~\\citep{Ling2009}. The nature and the mass of the compact object are still a subject for debate. Published results suggest either a neutron star of 1.4~$M_{\\odot}$~\\citep{Stark2003} or a black hole of less than 10~$M_{\\odot}$~\\citep{Hanson2000}. The identification of its donor star as a Wolf--Rayet star~\\citep{Vankerkwijk1992} classifies it as a high-mass X-ray binary. Nevertheless \\cyg\\ shows a short orbital period of 4.8~hr, typical of the low-mass binaries, which has been inferred from the modulation of both the X-ray~\\citep{Parsignault1972} and infrared emissions~\\citep{Becklin1973}.  Despite the strong X-ray absorption in this system, which may be caused by the wind of the companion star, the X-ray spectrum has been intensively studied. The source shows two main spectral X-ray states resembling the canonical states of the black hole binaries: the hard state (HS) and the soft state (SS; Zdziarski $\\&$ Gierlinski 2004; Hjalmarsdotter et al. 2008).  The HS is characterized by a weak soft thermal component and a strong non-thermal power-law emission peaking at hard X-ray energies, whereas the SS, though showing a non-thermal tail~\\citep{Szostek2008}, is dominated by the optically thick thermal disk emission. In Cygnus~X-3, however, the HS displays a high-energy cutoff at  $\\approx$ 20 keV, significantly lower than the $\\approx$ 100 keV value found for black hole binaries~\\citep[]{Hjalmarsdotter2004, Zdziarski2004}.  Adding to its peculiarity, \\cyg\\ is the brightest radio source among the X-ray binaries, quite frequently exhibiting huge radio flares~\\citep[as intense as few thousand times the quiescent emission level of $\\sim 20$ mJy at 1.5 GHz;][]{Braes1972} first seen by \\citet{Gregory1972}. During these outbursts, which occur mainly when the source is in the SS, and last several days, \\cyg\\ reveals the presence of collimated relativistic jets~\\citep[e.g.,\\ ][]{Marti2001,Mioduszewski2001,Geldzahler1983, Miller-Jones2004}. Thus, \\cyg\\ has been classified as a microquasar. Based on arcsecond-scale radio observations and their relation with soft X-ray emission, six X-ray/radio states have been identified: quiescent (flux densities $\\sim$ 100 mJy), minor flaring ( $<$ 1Jy), suppressed ($<$ 100 mJy), quenched ($<$ 30 mJy), major flaring ($>$ 1 Jy), and post-flaring~\\citep[]{Szostek2008,Tudose2009, Koljonen2010}. In the transition from the X-ray hard/radio quiescent state to the SS, the radio emission is strongly suppressed, and if it reaches the quenched level, the source usually produces a major radio flare~\\citep{Waltman94}.  \\cyg\\ has also historically drawn a great deal of attention due to numerous claims of detection at TeV and PeV $\\gamma$ rays, using early-days detectors in these energy ranges~\\citep[]{Vladimirsky1973, Danaher1981,Lamb1982,Dowthwaite1983,Samorski1983,Cawley1985, Chadwick1985,Bhat1986}. However, a critical analysis of these observations raised doubts about their validity~\\citep{Chardin1989}, and in recent years the more sensitive instruments have not confirmed those claims for energies above 500 GeV~\\citep[]{Shilling2001,Albert2008a}. Nevertheless, microquasars are believed to produce a very high energy (VHE) emission inside the jets: high density and magnetic fields provided by the accretion disk and by the companion star create favorable conditions for effective production of $\\gamma$ rays~\\citep[] {Levinson1996,Romero2003,Bosch-Ramon2006}. This radiation could have either an episodic nature due to the ejection of strong radio-emitting blobs~\\citep{Atoyan1999}, generally occurring in the SS in the case of \\cyg\\, or a (quasi) stationary character being originated in the persistent compact jet present during the HS~\\citep[]{Bosch-Ramon2006}.  Using data from the Energetic Gamma-Ray Experiment Telescope detector aboard the Compton Gamma-Ray Observatory, \\citet{Mori1997} reported an average flux of $(8.2 \\pm 0.9) \\times 10^{-7}$~photons cm$^{-2}$ s$^{-1}$ at energies above 100 MeV coming from the direction of \\cyg. However, no orbital modulation was detected in the signal, precluding a solid association. The experimental situation in the high-energy region has been drastically changed by the results recently published by \\emph{AGILE}~\\citep{Tavani2009} and \\emph{Fermi}/LAT~\\citep{Abdo2009}. \\emph{AGILE} detected the source in five different moments, for a couple of days each, four of them corresponding to the peak emissions shown in the detailed \\emph{Fermi}/LAT light curve (see below). The last detection, occurred in 2009 July, has not been published yet by the \\emph{AGILE} Collaboration (A. Bulgarelli et al., in preparation). On the other hand, \\emph{Fermi}/LAT detected a clear signal from Cygnus X-3 above 100 MeV during two periods of enhanced activity lasting for several weeks and coinciding with the source being in the SS. The measured flux is variable and shows an orbital modulation, which confirms the origin of the signal from the microquasar. The \\emph{AGILE} flux level is comparable with the one of the \\emph{Fermi}/LAT flux peaks, as high as 2.0~$\\times$~10$^{-6}$~photons~cm$^{-2}$ s$^{-1}$ above 100 MeV.  Observations of binary systems with imaging atmospheric Cherenkov telescopes (IACTs) in the VHE band have proven very fruitful in recent years, with the detection of the orbitally modulated $\\gamma$-ray emitters PSR\\,B1259-63~\\citep{Aharonian2005a}, LS\\,5039~\\citep{Aharonian2005b,Aharonian2006}, and LS~I~+61~303~\\citep[]{Albert2006,Albert2008b,Albert2009,Acciari2008, Acciari2009,Anderhub2009}. However, these systems may be different from Cygnus~X-3 since the radio and high-energy radiation could be produced in all three sources by the interaction of the winds of a star and an orbiting pulsar~\\citep[]{Maraschi1981,Tavani1997,Dubus2006}. Nevertheless, \\cyg\\ GeV $\\gamma$-ray modulation presents some common features to those observed in LS\\,5039 and LS~I~+61~303. On the other hand, although observations of other well-established microquasars, such as GRS~1915$+$105~\\citep[]{Acero2009,Saito2009}, did not reveal any signal, there is an interesting possibility of VHE emission from Cygnus~X-1~\\citep{Albert2007b}, still to be confirmed by an independent detection. The origin of this possible emission has been speculated to be linked to a maximum of the X-ray super-orbital modulation of the system~\\citep{Rico2008}. Moreover, there is also a claim by \\emph{AGILE} about a Cygnus~X-1 detection during a short flare (Sabatini et al. 2010, ATel 2512), which, however, has not been corroborated by \\emph{Fermi}/LAT yet\\footnote{Preliminary results can be found~at~\\emph{ http://fermisky.vlogspot.com/2010/03/lat-limit-on-cyg-x-1-during-reported.html}}. Such experimental results endorsed by the theoretical predictions, have encouraged deeper observations of the black hole binaries in the VHE band.  This paper reports observations of \\cyg\\ performed with the Major Atmospheric Gamma Imaging Cherenkov (MAGIC) telescope between 2006 and 2009. All these observations were carried out with the first stand-alone MAGIC telescope, MAGIC phase I. \\cyg\\ observations were planned to cover different X-ray/radio states, including those that showed a strong high-energy $\\gamma$-ray flux. This allowed us to search for VHE emission, above 250 GeV, from \\cyg\\ in the different cases for which this radiation is predicted in theoretical scenarios. In addition, specific analyses were performed to look for predicted features, such as periodic and variable emission. Section~\\ref{sec::obs}  describes the performance of MAGIC and the observational strategy. The analysis chain is explained in Section~\\ref{sec::data}. The results obtained by MAGIC in a multi-wavelength context are reported in Section~\\ref{sec::results}. A brief discussion is given in Section~\\ref{sec::discussion}. ",
        "conclusions": "\\label{sec::discussion}  MAGIC observations of \\cyg\\ cover all X-ray spectral states of the source in which VHE emission is thought to be likely produced either from a persistent jet in the HS or during powerful ejections in the SS. However, no significant excess events were found in any of the inspected samples.  VHE $\\gamma$ rays have been predicted from microquasar jets~\\citep[]{Atoyan1999,Romero2003,Bosch-Ramon2006}. A robust prediction of modeling is that photon--photon absorption cannot be neglected if the $\\gamma$ rays are produced close to a massive star~\\citep[]{Bednarek1997,Orellana2007}. In particular for \\cyg\\, the presence of a Wolf--Rayet companion, with temperature $T_*\\approx 10^5$~K and radius $R_*\\approx 10^{11}$~cm, leads to an optical depth $\\tau\\ge 1$ for VHE $\\gamma$ rays for an emitter located at several orbital radii from the star~\\citep{Bednarek2010}. Even under very efficient electromagnetic cascading, i.e., a radiation to magnetic energy density ratio $8\\,\\pi\\,u_*/B^2\\gg 1$, the expected VHE fluxes are below the sensitivity of the present instruments~\\citep{Bednarek2010}. Therefore, in order to detect VHE photons, the emitter should be located far from the binary system.  \\emph{Fermi}/LAT detected \\cyg\\ when it was in the SS and found orbital modulation for the radiation above $100$~MeV with a photon index of 2.7 for the periods of enhanced activity~\\citep{Abdo2009}. For the epochs outside these high-activity periods, the GeV flux decreases significantly (Figure~\\ref{fig::multi} in~\\citet{Abdo2009}) and no modulation is found. The GeV orbital light curve of \\cyg\\ in the high-energy active periods can be explained in the context of anisotropic inverse Compton scattering with the stellar photons~\\citep[]{Abdo2009,Bosch-Ramon2009,Dubus2010}, which is also energetically more efficient than hadronic mechanisms such as $pp$ interactions or photomeson production. Since very bright X-ray emission is produced in the inner accretion disk or at the base of the jets in Cygnus~X-3, the GeV radiation would be absorbed unless it is originated beyond $\\sim 10^{10}$~cm above the compact object. This implies that the GeV radiation is produced in the jet of \\cyg\\ rather than in the inner accretion disk/corona region. On the other hand, the GeV emitter cannot be too high in the jet, since otherwise there would not be strong orbital modulation (see also Abdo et al. 2009). Therefore, the observed GeV and the predicted detectable VHE emission cannot be explained by one particular population because the former should be produced in/close to the system and the latter farther from it. The location of a hypothetical VHE emitter could coincide with the innermost region of the radio emitting jet, which, to avoid synchrotron self-absorption, should start relatively far from the binary system. This is consistent with the fact that MAGIC did not detect \\cyg\\ during the high-activity GeV period, even though the flux ULs are close to a power-law extrapolation of the \\emph{Fermi}/LAT spectrum to energies greater than 100 GeV and well below an extrapolation of the \\emph{AGILE} spectrum.  During the periods of GeV high-activity of the source, as by the synchrotron self Compton scenario~\\citep{Atoyan1999}, a detectable TeV signal could arise during the first hours of a radio outburst. Unfortunately, MAGIC has never observed the source during this phase of the flare, but always some days before or after the maximum radio flux. This radiation would not be strongly modulated, due to its origin far from the system. The two times more sensitive two telescopes arrangement, MAGIC phase II, may indeed detect \\cyg\\ if it observes the source for longer time at the very maximum of a GeV flare, which might be even earlier than the onset of the radio outburst~\\citep{Abdo2009}.  In the HS, the VHE emission is expected to be produced inside the compact and persistent jets, whose total luminosity is estimated to be at least 10$^{37}$~erg~s$^{-1}$~\\citep{Marti2005}. The MAGIC VHE $\\gamma$-ray UL set at 1.1~$\\times$~10$^{-12}$~erg~cm$^{-2}$~s$^{-1}$ is equivalent to a VHE luminosity of $\\simeq$ 7~$\\times$~10$^{33}$~erg~s$^{-1}$ at 7 kpc. Thus, the maximum conversion efficiency of the jet power into VHE $\\gamma$ rays is 0.07$\\%$ which is similar to that of Cygnus~X-1 for the UL on the VHE steady emission, but 1 order of magnitude greater than that of GRS~1915+105~\\citep[]{Albert2007b, Acero2009}. These ULs are in good agreement with the theoretical expectations which generally predict a VHE steady luminosity of $\\simeq$ 10$^{32}$~erg~s$^{-1}$. Persistent galactic jets do not seem to be good candidate sources to be detected at VHEs by the current sensitivity instrumentation. Only 10 times more sensitive future instruments, such as Cherenkov Telescope Array, may have a chance to detect such VHE emission. This would provide a new handle on the emission mechanisms of compact jets."
    },
    "1005/1005.5165_arXiv.txt": {
        "abstract": "We use high-resolution three-dimensional adaptive mesh refinement simulations to investigate the interaction of high-redshift galaxy outflows with low-mass virialized clouds of primordial composition. While atomic cooling allows star formation in objects with virial temperatures above $10^4$ K, ``minihaloes\" below this threshold are generally unable to form stars by themselves. However, these objects are highly susceptible to triggered star formation, induced by outflows from neighboring high-redshift starburst galaxies. Here we conduct a study of these interactions, focusing on cooling through non-equilibrium molecular hydrogen (H$_2$) and hydrogen deuteride (HD) formation. Tracking the non-equilibrium chemistry and cooling of 14 species and including the presence of a dissociating background, we show that shock interactions can transform minihaloes into extremely compact clusters of coeval stars.  Furthermore, these clusters are all less than $\\approx 10^6 M_\\odot,$ and they are ejected from their parent dark matter halos: properties that are remarkably similar to those of the old population of globular clusters. ",
        "introduction": "A generic prediction of the cold dark matter model is a large high-redshift population of gravitationally-bound clouds  that are unable to form stars.  Because atomic H and He line cooling is only effective at temperatures above $10^4$ K, clouds of gas and dark matter  with virial temperatures below this threshold must radiate energy through dust and molecular line emission.   While the levels of H$_2$ left over from recombination are sufficient to cool gas in the earliest structures (\\eg Abel \\etal 2002; Bromm \\etal 2002), the resulting 11.20 - 13.6 eV background emission from the stars in these objects (\\eg Haiman \\etal 1997; Ciardi \\etal 2000; Sokasian \\etal 2004; O'Shea \\& Norman 2007) is likely to have quickly dissociated these trace levels of primordial molecules (Galli \\& Palla 1998).   And although an early X-ray background could have provided enough free electrons to promote H$_2$ formation, the relative strength between these two backgrounds is uncertain, and it is unlikely that the background was strong enough to balance ultraviolet (UV) photodissociation. Even if there were some trace amount of H$_2$ in these clouds, it is likely to be in such a small abundance as to not impact their structure (Whalen \\etal 2008a; Ahn \\etal 2009).  The result is a large number of dark matter halos that were massive enough to overcome the thermal pressure of the primordial intergalactic medium and retain their gas,  but not massive enough to excite the radiative transitions necessary to cool this gas into stars.  These ``minihalos,'' whose virial temperatures were $T_{\\rm{vir}} \\lesssim 10^{4}$ K and whose total masses were between $10^{4}$ and $10^{7.5} M_{\\odot}$ would have remained sterile until some outside force either disrupted them, or more interestingly, disturbed them so as to catalyze coolant formation.  In fact, two possible coolant formation methods have been considered in detail in the literature: ionization fronts and shock fronts.  In the case of ionization fronts, such as would occur during the epoch of reionization, high-energy photons emitted from galaxies or quasars interact with the neutral atomic minihalo gas.  Bond, Szalay \\& Silk (1988) originally  discussed how the resulting photoionization  would expel the gas contained in a minihalo by suddenly heating it to $T \\approx  10^4$ K, as would be the case in the optically-thin limit.  On the other hand,  Cen (2001), used simple analytic estimates to argue that ionization fronts would cause non-equilibrium H$_2$ formation and the collapse of the gas inside the gravitational potential.  However, Barkana \\& Loeb (1999) studied minihalo evaporation using static models of uniformly illuminated spherical clouds, accounting for optical depth and self-shielding effects, and showed that the cosmic UV background boiled most of the gas out of these objects.  Later, Haiman, Abel, \\& Madau (2001) carried out three-dimensional (3D) hydrodynamic simulations  assuming the minihalo  gas was spontaneously heated to $10^4$ K,  also finding quick disruption. Finally, full radiation-hydrodynamical simulations of ionization front-minihalo interactions were carried out in Iliev \\etal (2005) and Shapiro \\etal (2004;  see also Shapiro, Raga \\& Mellema 1997; 1998). These demonstrated that intergalactic ionization fronts decelerated when they encountered the dense, neutral gas inside minihaloes and were thereby transformed into D-type fronts, preceded by shocks that completely photoevaporated the minihalo gas.  A second and more promising avenue for coolant formation is the interaction between galactic outflows and minihaloes. These galaxy-scale winds, which are driven by core-collapse supernova and winds from massive stars, are commonly observed around dwarf and massive starbursting galaxies at both low and high redshifts  (\\eg Lehnert \\& Heckman 1996; Franx \\etal 1997; Pettini \\etal 1998; Martin 1999; 1998; Heckman \\etal 2000; Veilleux \\etal 2005; Rupke \\etal 2005), and a variety of theoretical arguments suggest that these galaxies represent only the tail end of a larger population of smaller  ``pre-galactic,\" starbursts that formed before reionization (Scannapieco, Ferrara \\& Madau 2002; Thacker, Scannapieco, \\& Davis 2002). Furthermore, the interstellar gas swept up in a starburst-driven wind can effectively  trap the ionizing photons behind it (Fujita \\etal 2003), meaning that at high redshifts, many intergalactic regions may have been impacted by outflows well before they were ionized.  Such $\\approx 100-300$ km/s shocks can cause intense cooling through two mechanisms: (i) the mixing of metals with ionization potentials below 13.6 eV (Dalgarno \\& McCray 1972), which allow for atomic line cooling even at temperatures below $10^4$ K; and (ii) the formation of H$_2$ and HD by nonequilibrium processes (Mac Low \\& Shull 1986; Shapiro \\& Kang 1987; Kang \\etal 1990; Ferrara 1998; Uehara \\& Inutsuka 2000), which allow for molecular line cooling associated vibrational and rotational transitions (Palla \\& Zinnecker 1988).  In fact, Scannapieco \\etal (2004) showed that these effects were so large that shock interactions could induce intense bursts of cooling and collapse in previously ``sterile'' minihalo gas. Using simple analytic models, they found that the most likely result was the formation of compact clusters of coeval stars, although they also emphasized the importance of multidimensional numerical studies to confirm this result.  In this paper, we present the first in a series of 3D numerical studies to capture  the physical interactions between primordial minihalos and supernova-driven galactic outflows.  While triggered star formation has been simulated in low-redshift intergalactic clouds impacted by radio jets (\\eg van Breugel \\etal 1985; Fragile \\etal 2004; Wiita 2004; Klamer \\etal 20004), this has never been numerically simulated in the high-redshift, minihalo case.   In this paper, we  focus on the physics of non-equilibrium H$_2$ and HD chemistry and associated cooling in determining the properties of the post-shock minihalo gas.  In particular, we show that molecule formation and the resulting cooling is strong enough to induce rapid minihalo collapse and star formation, leading to compact stellar clusters.  The structure of this paper is as follows.  In \\S 2 we outline the physical components of the model, focusing on our primordial H$_2$ and HD chemical network and cooling routines and their respective tests.  In \\S 3 we outline the general model used for the galactic outflow and the minihalo, and in \\S 4 we present the simulation results and relate the resulting stellar clusters to local observations. Conclusions are given in \\S 5. ",
        "conclusions": "Cosmological minihalos provide the initial building blocks out of which larger dark mater halos and galaxies form.  During reionization, these minihalos surrounded  all young galaxies and consisted of metal-free neutral atomic gas. The lack of molecules in these objects stalled their evolution at virial temperature below $\\approx 10^4 \\rm \\ K$ as cooling from atomic hydrogen becomes inefficient and any molecular coolants are dissociated by UV Lyman-Werner photons, too hot for star formation in such small objects. Only interactions that cause non-equilibrium molecule formation could allow the gas in minihalos to collapse and form stars.  Previous work has used ionization fronts to create these conditions (Cen 2001), however, 3D hydrodynamic simulations show that for either a stellar or quasar source, instead of creating clouds of H$_2$, the cloud is instead completely photo-evaporated (Iliev \\etal 2005; Shapiro \\etal 2006). However galactic outflows are another option for triggering star formation. Shocks provide the conditions for non-equilibrium chemistry without completely destroying the cloud.  To simulate these interactions, we have added a primordial chemistry network and associated cooling routines into FLASH3.1. This network traces collisional ionization and recombination of hydrogen and helium as well as the formation of two primary coolants in the absence of metals: H$_2$ and HD. We have also included the impact of a dissociating background on these rates and implemented routines that control the cooling of gas  in the presence atomic and molecular line cooling.  With these code developments in place, we  have been able to simulate the interaction between a galactic outflow and a primordial minihalo that results in structures that are identified as proto-globular clusters. The shock fills two very important roles. First, it ionizes the neutral gas found in the minihalo, which recombines and begins to form  H$_2$ and HD, through nonequillbrium processes. These coolants allow the cloud to cool to much lower temperatures, triggering star formation. Secondly, the shock imparts momentum into the gas and accelerates it above the escape velocity. This creates a cloud of dense, cold molecular gas that is free from dark matter halos.  Together, these many similarities suggest that there is not only room at low redshifts for the kinds of clusters were see in our simulations, but that low-redshift observations {\\em require} the formation of stars by a mechanism remarkably similar to the one we are simulating.   This physical connection requires further investigation, especially in light of a  final remarkable property of globular clusters, the chemical homogeneity of their stars.  While there are large metallicity between different globular clusters the dispersion of [Fe/H] within  any given cluster  is less than 0.1 dex, or a factor of $\\approx$ 1.25 (Suntzeff 1993).  Usually, this is explained either by  pre-enrichment,   meaning that stars form out of gas that had already been homogeneously enriched by a previous generation of supernovae (\\eg Elmegreen \\& Efremov 1997; Bromm \\& Clarke 2002),  or ``self-enrichment,\" meaning that the protocluster cloud was enriched by one or more SN events occurring within it (e.g., Brown \\etal 1995; Cen 2001; Nakasato \\etal 2000; Beasley et al. 2003; Li \\& Burstein 2003).   Both types of scenarios have strong flaws.  In the pre-enrichment case,  the key problem is just exactly what the previous population was, why it played only a secondary role in the formation history of the cluster, and how it could have enriched this material on very short timescales. In the self-enrichment picture, on the other hand, the key problems is that the kinetic energy corresponding to resulting SN is likely to unbind the gaseous protocluster (Peng \\& Weisheit 1991) long before it enriches the gas.  Recent two dimensional numerical simulations of primordial supernova inside minihalos  show that they do in fact unbind these systems over a range of supernova types and halo sizes (Whalen \\etal 2008b).  Shock minihalo interactions have the potential to combine the best features of both these scenarios, bringing in the metals from an outside source, such that the gas collapses instead of unbinds, but nevertheless explaining how star formation could be synchronized in a large mass of coolant-filled gas.  Yet highly homogenous population of stars can only be formed if the incoming, enriched gas and the primordial, minihalo material are sufficiently well mixed before star formation begins.   Modeling this process will require simulations that build on the ones here to include metal line cooling above and below $10^4$ K and detailed models of subgrid-mixing.   Testing this assumption and determining what sets of parameters lead to realistic globular cluster formation will be the focus of the upcoming papers in this series."
    },
    "1005/1005.2868_arXiv.txt": {
        "abstract": "We revise and extend the extreme value statistic, introduced in \\cite{gup08}, to study directional dependence in the high redshift supernova data; arising either from departures from the cosmological principle or due to direction dependent statistical systematics in the data. We introduce a likelihood function that analytically marginalises over the Hubble constant, and use it to extend our previous statistic. We also introduce a new statistic that is sensitive to direction dependence arising from living off-centre inside a large void, as well as previously mentioned reasons for anisotropy. We show that for large data sets this statistic has a limiting form that can be computed analytically.  We apply our statistics to the gold data sets from \\cite{rie04} and \\cite{rie07}, as in our previous work. Our revision and extension of previous statistic shows that 1) the effect of marginalsing over Hubble constant instead of using its best fit value has only a marginal effect on our results. However, correction of errors in our previous work reduce the level of non-Gaussianity in the 2004 gold data that was found in our earlier work.  The revised results for the 2007 gold data show that the data is consistent with isotropy and Gaussianity. Our second statistic confirms these results. ",
        "introduction": "Observations of high-redshift supernovae of Type~Ia \\citep{per99,rie98,rie02,rie04}, along with observations of the cosmic microwave background \\citep{ben02,pag06}, indicate a flat universe with an accelerating expansion. The existing evidence is consistent with a universe that is dominated today by a cosmological constant term in the Einstein's theory of gravity, however, due to fine tuning required in this model, several alternative explanations have been suggested \\citep{silv09,frie08,sahn06}. If we phenomenologically treat this unknown energy component as an ideal fluid --- called the dark energy --- then the pressure of this fluid has to be negative, that is, for an equation of state $p=w\\rho$, we have $\\rho > 0$ and $w<0$. In this description the cosmological constant model has an equation of state given by $p=-\\rho$, and thus we expect that a phenomenological modeling of the cosmological data with a constant $w$ should give $w \\simeq -1$, which turns out to be the case; the current data indicating one-third of the total density in the form of dark and baryonic matter and two-thirds in the form of the cosmological constant --- a model that is known as the $\\Lambda$CDM model. However, it should be noted that the equation of state for the dark energy could differ from $w=-1$ in the past when the dark energy is expected to be subdominant (for consistency with constraints arising from the observations of the microwave background). Indeed, in many plausible models of dark energy the equation of state approaches $w=-1$ only in the recent past and thus may show only tiny departures from $w=-1$ at low redshifts, while at high redshifts it might remain subdominant and therefore would have a weaker observational signature. Although it is perhaps hard to extend supernova (SN) observations beyond $z \\simeq 3$ or so, neutral hydrogen observations of the post-reionization epoch may be able to provide constraints on dark energy all the way up to $z \\simeq 6$ \\citep{bhar09}.  To resolve this issue we require data that is precise enough to discern tiny variations in the dark energy. It is also required that data be available at a large number of redshifts to constrain the detailed temporal behaviour of dark energy. At present the only data that comes reasonably close to these requirements is provided by the observations of the high redshift type Ia supernovae (SNe), which are believed to be standard candles --- a claim that may be doubtful considering the relatively poorly understood physics, and the possibility of physical mechanisms such as dust in the inter-galactic medium that systematically dims them. Another concern is that the SN data sets are usually collated from several different experiments that might have slightly different systematics, due either to instrumental effects or the fact that they observe different directions in the sky; for example systematic errors in correcting Galactic or source galaxy dust extinction might leave residual anisotropies in data. These considerations imply that to obtain precise information about the behaviour of dark energy we should first have a good knowledge of the statistical properties of SNe, both random as well as systematic.  The standard model of cosmology is based on the Cosmological Principle \\citep{pee93}, which states that the Universe is homogeneous and isotropic on large scales.  Recent work provides some evidence for what is known as the Hubble Bubble \\citep{zeh98,jha07} indicating that we might be living inside a large void that has a different value of the Hubble constant inside from what is outside the bubble. There is evidence for such large scale voids in the CMB maps as well \\citep{cruz05,olv06,cruz06,cruz07}, suggesting that such large voids are not implausible. There is no reason for us to believe that we are living at the precise centre of such a void. Therefore, if we happen to be living in such a void then various distance measures, such as the luminosity distance and the angular diameter distance, may not be isotropic. \\cite{gup08} (hereafter GSL08) used the extreme value statistics to show that the two SN data sets, \\cite{rie04} (GD04) and \\cite{rie07} (GD07), do show some evidence for direction dependence. Several other works have also indicated either systematic problems with the high-redshift SN data or directional dependence in the SN data and other probes \\citep{kol01,ness04,ness05,ness07, jain106,jain206,coo09}.  We recently found that our analysis in GSL08 contained mistakes due to a coding error.  In this paper our main task is 1) to provide revised results and to extend the previous work by introducing a statistic that marginalises over the Hubble parameter instead of using its best fit value 2) to introduce a new statistic that has a greater sensitivity to the signatures of living off-centre inside a large void. We also give asymptotic form for this statistic, which makes it easier to use it for large data sets.  The plan of the paper is as follows. In \\S~2 we introduce the likelihood function marginalised over the Hubble constant and provide a comparison with what we obtain using the Bayesian statistic. In \\S~3 we introduce the two statistics. In \\S~4 we present our results and \\S~5 we end with conclusions. ",
        "conclusions": "\\begin{figure} \\centering \\includegraphics[width=0.50\\textwidth]{figure5.ps} \\caption{The theoretical, analytic and the bootstrap probability distributions for the $\\Delta_{\\chi}$ statistic for GD04. The analytic distribution uses the limiting values of the shape and position parameters for the Gumbel distribution. Since the total number of independent directions is most likely smaller than $N_d$, we overestimate the position of peak. The bootstrap distribution indicates anisotropy at the level of slightly greater than about two sigma.  } \\label{fig:new04} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=0.50\\textwidth]{figure6.ps} \\caption{The theoretical, analytic and the bootstrap probability distributions for the $\\Delta_{\\chi}$ statistic for GD07 as in Figure~\\ref{fig:new04}. The bootstrap distribution indicates consistency with isotropic distribution of SNe in GD07. } \\label{fig:new07} \\end{figure} We have presented corrected results for the $\\Delta$ statistic introduced in GSL08 and showed that the results for the $\\Delta_{\\chi^2}$ statistic change substantially after correcting for a numerical bug in the our code. Our present method extends the previous work by invoking an estimator that is marginalised over the Hubble parameter. However, our results are not affected substantially due to this change since the bootstrap distribution and the location of $\\Delta_{\\chi^2}$ does not change substantially for GD04 compared to our previous results. However, for GD07 we have corrected another mistake in our earlier analysis and we find that the bootstrap distribution as well as the position of GD07 with respect to it have changed.  Our main conclusions are 1) the match between the bootstrap and the theoretical distribution is much better than was presented in GSL08 2) GD04 shows some evidence for non-Gaussianity, however, GD07 is statistically consistent with a Gaussian distribution of residuals 3) and although GD04 has a weak direction dependence, GD07 is consistent with an isotropic distribution of SNe.  Our new statistic $\\Delta_{\\chi}$ similarly shows a weak direction dependence in GD04 but no significant anisotropy in GD07. One surprising feature we find is that for both GD04 and GD07, $\\Delta_{\\chi^2}$ turns to be larger than the mode of the distribution, while $\\Delta_{\\chi}$ is smaller than the mode of the distribution. However, since our results are consistent with isotropy, we do not investigate the implications of this puzzling feature. The results of this statistic can be compared to the results of \\cite{coo09}. The directions of maximum anisotropy are different. However, since the methods and data are different, and the results suggest only a weak direction dependence, this mismatch is not unexpected.  \\begin{table} \\begin{center} \\caption{ The model parameters for GD04 and GD07 are tabulated  here.} \\label{tbl-newfit} \\bigskip  \\begin{tabular}{crrrrrr} \\hline Model & Set & $\\Om$ & $ H_0 $ &  $\\chi^2$ \\\\ \\hline $\\Lambda$CDM & GD04 & 0.30 & 64.5 &  177.3 \\\\ $\\Lambda$CDM & GD07 & 0.33 & 63.0 &  158.8 \\\\ \\hline\\hline \\end{tabular} \\end{center} \\end{table}    \\begin{table} \\begin{center} \\caption{ Direction for maximum $\\Delta$ GD04 and GD07 are tabulated  here.} \\label{tbl-newdel} \\bigskip  \\begin{tabular}{crrrrr} \\hline Model & Set & $\\Delta_{\\chi}$ &  long  &  lat \\\\ \\hline $\\Lambda$CDM & GD04 &-1.52 & 107 &  26 \\\\ $\\Lambda$CDM & GD07  &-1.77  & 97 &  70 \\\\ \\hline\\hline \\end{tabular} \\end{center} \\end{table}    \\appendix \\label{appendix}"
    },
    "1005/1005.4952_arXiv.txt": {
        "abstract": "We use a hybrid test particle/N-body simulation to integrate 4 million massless test particle trajectories within a fully self-consistent $10^5$ particle N-body simulation. The number of massless particles allows us to resolve fine structure in the spatial distribution and phase space of a dwarf galaxy as it is disrupted in the tidal field of a Milky Way type galaxy. The tidal tails exhibit nearly periodic clumping or a smoke-like appearance. By running simulations with different satellite particle mass, halo particle mass, number of massive and massless particles and with and without a galaxy disk,  we have determined that the instabilities are not due to numerical noise, amplification of structure in the halo, or shocking as the satellite passes through the disk of the Galaxy.  We measure Jeans wavelengths and growth timescales in the tidal tail and show that the Jeans instability is a viable explanation for the clumps. We find that the instability causes velocity perturbations of order 10 km/s.  Clumps in tidal tails present in the Milky Way could be seen in stellar radial velocity surveys as well as number counts. We find that the unstable wavelength growth is sensitive to the simulated mass of dark matter halo particles. A simulation with a smoother halo exhibits colder and thinner tidal tails with more closely spaced clumps than a simulation with more massive dark matter halo particles. Heating by the halo particles increases the Jeans wavelength in the tidal tail affecting substructure  development, suggesting an intricate connection between tidal tails and dark matter halo substructure. ",
        "introduction": "There is evidence for past and ongoing accretion of small objects by the Milky Way halo, the most dramatic object being the disrupted Sagittarius dwarf galaxy \\citep{ibata94}. Disrupting satellites leave behind tidal tails and, on longer timescales, stellar streams \\citep{bekki03,helmi03,meza05,penarrubia05,purcell07,helmi08}. Previous studies of stellar streams have used N-body simulations to study the disruption and evolution of merging galaxies (e.g., \\citealt{helmi06,johnston08,gomez10}). However, most previous simulations have not placed many particles in the disrupting object itself.  A large number of particles is required to resolve structure in the velocity distribution in a small local volume (e.g., \\citealt{minchev09,gomez10}). After satellite disruption, satellite particles are distributed all over the Galaxy.  Early simulations of dwarf galaxy disruption necessarily contained few particles. The focus of many of these simulations was constraining the orientation and shape of the halo from observations of the Sagittarius Dwarf stream rather than searching for substructure in the tails themselves. The simulations by \\citet{johnston96,johnston98} contained only $10^4$ particles in the dwarf galaxy and those by \\citet{helmi06} only had 5000 particles. The simulations by \\citet{bullock05} had $10^4$ massive particles but an additional $1.2 \\times 10^5$ massless test particles in their dwarf satellites. The simulations discussed by \\citet{gomez10} are described in more detail by \\citet{villalobos08} and contained $3 \\times 10^5$ particles in the satellite.   While substructure in the form of clumps has not been detected in the Sag Dwarf streams, it has been seen in tidal tails associated with smaller objects. Clumped structure observed in the tidal debris of the globular cluster Palomar 5 \\citep{oden02} is interpreted in terms of oscillations in the cluster \\citep{gnedin99} caused by a previous passage through the Galactic disk \\citep{oden02,dehnen04}, epicyclic perturbations excited during tidal disruption \\citep{kupper08} and Jeans instability \\citep{quillen10}. Alternative possibilities accounting for structure in cluster tidal tails include the effect of dark matter subhalos, as explored by \\citet{ibata02,mayer02,johnston02,penarrubia06,valluri08,carlberg09}.  Here, we strive to carry out N-body simulations with a larger number of particles placed in the disrupting dwarf galaxy. We use the simulations to look in detail at the structure of the  tidal tails during and following disruption from the dwarf galaxy. ",
        "conclusions": "We have used hybrid test particle/N-body integrations to increase the number of particles integrated within an N-body integrator so that we can more accurately resolve substructure. We have used our code to study the tidal disruption of a dwarf galaxy in a polar orbit of a Milky Way type galaxy.  We have placed additional test particles in the dwarf galaxy so that we can resolve fine structure in the tidal tails.  In our simulations, we have found that stellar tidal tails can exhibit periodic ridges oriented perpendicular to the orbit.  Such structure had not previously been noticed in similar N-body simulations, possibly because the tidal tails did not contain sufficient numbers of particles. However similar structure has been seen in simulations of tidal tails that include gas \\citep{wetzstein07} and in this setting the clumps are interpreted in terms of a gravitational instability.  By comparing simulations with different numbers of massive dwarf and halo particles and with and without a disk, we have considered several explanations for the formation of clumps. We have ruled out the following methods for the formation of periodically spaced over-densities in our simulation: shocking by the galactic disk, numerical noise associated with under-populating the dwarf galaxy, amplification of structure in the halo, and epicyclic motions in the tidal tails.  We have measured the Jeans wavelength prior to the growth of the substructure and found that the tails are unstable to Jeans instability. The wavelengths of subsequently formed clumps are approximately consistent with the Jeans wavelength measured in the tail prior to formation. The timescale for growth is approximately consistent with the estimated growth timescale.  These estimates suggest that Jeans instability is a viable interpretation of the clumps exhibited by our simulations. We find that the spacing between clumps is sensitive to the mass of our simulated dark matter halo particles.  This is likely because heating by dark matter particles can increase the Jeans wavelength of the tidal tail.  We find that the clumps are also visible in a radial velocity projections suggesting that Jeans instabilities may be observable in tidal tails in our galaxy not only number counts but also in phase space using comparisons of radial velocity versus distance or position on the sky.  In the future, we expect increasingly rich data samples expanding the number of stars in the Milky Way with measured properties.  These surveys may make it possible to probe for or rule out substructure in tidal tails such as exhibited by our simulations.   If Jeans instabilities occur in tidal tails then associated heating caused by them should not be interpreted in terms of heating by halo substructure alone. Furthermore the fastest growing unstable wavelength may be sensitive to heating from the dark matter substructure implying that there may be a complex connection between halo and tidal tail substructure.  The simulations carried out here were made using a modest direct N-body code on a desktop computer. Future studies could test the results presented here with a more sophisticated N-body code (such as a tree code), and integrate more particles by doing the simulations on a supercomputer. Further studies can be carried out in different mass and stripping regimes for the disrupting object, such as globular clusters. We would also be interested in carrying out simulations that would more fully probe the possible connection between dark matter substructure and tidal tail morphology by simulating a more detailed halo than considered here.  \\vskip 2.0truein We thank Larry Widrow for giving us and helping us with his code GalacticICS. We thank Jeff Bailin, Chris Purcell and Heidi Newberg for helpful communications. The King model was generated using code made available\\footnote{http://www.physics.mcmaster.ca/~syam/software.html} by Sergey Mashchenko. We thank Evghenii Gaburov and Stefan Harfst for making $\\phi$GRAPE and Sapporo available. Support for this work was provided by NSF through award AST-0907841."
    },
    "1005/1005.1939_arXiv.txt": {
        "abstract": "{Infrared-dark clouds (IRDCs) are the precursors to massive stars and stellar clusters.  G011.11$-$0.12 is a well-studied filamentary IRDC, though, to date, the absence of far-infrared data with sufficient spatial resolution has limited the understanding of the structure and star-formation activity.} {We use \\it Herschel \\rm to study the embedded population of young pre- and protostellar cores in this IRDC. } {We examine the cloud structure, which appears in absorption at short wavelength and in emission at longer wavelength.  We derive the properties of the massive cores from the spectral energy distributions of bright far-infrared point sources detected with the PACS instrument aboard \\it Herschel. \\rm} {We report on the detection and characterization of pre- and protostellar cores in a massive filamentary infrared-dark cloud G011.11$-$0.12 using PACS.  We characterize 18 cores directly associated with the filament, two of which have masses over 50~$\\msun$, making them the best candidates to become massive stars in G011.11$-$0.12.  These cores are likely at various stages of protostar formation, showing elevated temperature ($<$T$> \\sim$ 22~K) with respect to the ambient gas reservoir.  The core masses ($<$M$> \\sim$ 24~$\\msun$) are small compared to that in the cold filament.  The mean core separation is 0.9~pc, well in excess of the Jeans length in the filament.  } {We confirm that star formation in IRDCs is underway and diverse, and IRDCs have the capability of forming massive stars and clusters.} ",
        "introduction": "Star formation occurs in cold, dense molecular cores.  In local low-mass star-forming regions, this process takes place in relatively isolated environments and is better understood compared to the distant regions in which massive stars form.  Specifically, it is the identification of the cold precursors to stars - the so-called ``pre-stellar cores'' that has led to a greater understanding of the initial conditions of low-mass star formation;  however, the high-mass counterpart has been more difficult to identify.  Thus it is not clearly understood how the early stages of massive star formation compare to that of low-mass stars.  Infrared-dark clouds (IRDCs) are complexes of cold (T $<$ 20~K), dense ($n > 10^4$~cm$^{-3}$) molecular gas, now believed to be the precursors to massive stars and star clusters \\citep[e.g.][]{rathborne2006}, although there is diversity among this class of objects.  Early observations revealed that the cold dust in IRDCs absorbs the bright Galactic background emission at 8~$\\mu$m \\citep{egan_msx, hennebelle_isogal} and emits at sub-mm wavelengths \\citep[e.g.][]{carey_submmIRDC}.  Later, high-resolution studies found that IRDCs exhibit significant substructure \\citep[e.g.][]{ButlerTan2009, ragan_spitzer}, and (sub-)millimeter studies show that they host massive embedded star formation \\citep[e.g.][]{Rathborne_2007_protostars, Beuther_protostars_IRDC, Hennemann2009}.  The spectral energy distribution (SED) of massive embedded protostars is expected to peak in the far-infrared but is usually very faint in the mid-infrared.  The coldest cores are not readily detectable with {\\em Spitzer} because of sensitivity limitations at 24~$\\mu$m and confusion at longer wavelengths due to limited spatial resolution.  With {\\em Herschel} \\citep{A&ASpecialIssue-Herschel}, we now have the spatial resolution and the capability to perform large-area maps to search for these deeply embedded objects and to characterize the SEDs.  The large, filamentary IRDC G011.11$-$0.12 was selected as a target in the Early Phases of Star Formation (EPoS) {\\em Herschel} key program (PI: O. Krause).  Together with complementary {\\em Spitzer} data, we use PACS \\citep{A&ASpecialIssue-PACS} and SPIRE \\citep{A&ASpecialIssue-SPIRE} observations to constrain the properties of the massive embedded cores within this IRDC.  \\begin{figure*} \\centering \\includegraphics[scale=0.47]{f1.ps} \\caption{The IRDC G011.11$-$0.12 from {\\em Spitzer} to SCUBA wavelengths.  (a) {\\em Spitzer} IRAC 8$\\mu$m image.  (b) {\\em Spitzer} MIPS 24~$\\mu$m image.  The peak SCUBA positions (P1, P2,...) as assigned by \\citet{Johnstone_G11} are indicated.  (c) {\\em Herschel} PACS 70~$\\mu$m image.  (d) {\\em Herschel} PACS 100~$\\mu$m image with the positions of analyzed sources (see Sect. 4) indicated.  Sources with no 24~$\\mu$m counterpart are circled in cyan, sources with 24~$\\mu$m counterparts are circled in green.  (e) {\\em Herschel} PACS 160~$\\mu$m image.  (f) {\\em Herschel} SPIRE 250~$\\mu$m image.  (g) {\\em Herschel} SPIRE 350~$\\mu$m image.  (h) 8~$\\mu$m extinction map (derived from {\\em Spitzer} data) with (blue) SCUBA 850~$\\mu$m continuum contours (levels are 0.4 0.8 1.2 Jy beam$^{-1}$).  Source labeling same as in panel (d).  The approximate beam size is indicated in the lower-left corner of each panel.} \\label{fig:images} \\end{figure*}  Due to its high absorbing contrast against the Galactic background emission (particularly at 8~$\\mu$m), the IRDC G011.11$-$0.12 ($d$~=~3.6~kpc) has been subject to detailed studies.  \\citet{carey_msx} confirm the presence of dense molecular gas based on H$_2$CO observations.  \\citet{carey_submmIRDC} report 850 and 450~$\\mu$m continuum observations of the thermal dust emission which precisely follows the structure seen in absorption in the mid-infrared, and the column density reaches up to 10$^{23}$~cm$^{-2}$ \\citep[see also][]{Parsons2009}.  The so-called P1 position is seen in emission in both 8 and 850~$\\mu$m images (Fig. 1), a likely indicator of embedded star formation.  \\citet{Pillai_ammonia} observe G011.11$-$0.12 as part of an ammonia survey, including the (1,1) and (2,2) inversion transitions with 40$\\arcsec$ spatial resolution.  The masses of four sub-regions (P1, P2, P6, and P7 in the \\citet{Johnstone_G11} nomenclature) were found to be several hundred $\\msun$ each, though other mass-tracing techniques find higher values for the total filament mass \\citep[e.g.][]{Marshall2009}.  The NH$_3$ observations confirm that temperatures are very low throughout the filament, averaging 12~K.   In follow-up studies, observations of masers \\citep{Pillai_G11}, extended excess 4.5~$\\mu$m emission \\citep[a tracer of shocked H$_2$][]{Cyganowski2008}, and molecular line emission \\citep{Pillai_G11, Chen_infallEGO} have all supported the presence of ongoing star formation within the filament. ",
        "conclusions": ""
    },
    "1005/1005.3427_arXiv.txt": {
        "abstract": "The high-redshift gamma-ray bursts (GRBs), GRBs 080913 and 090423, challenge the conventional GRB progenitor models by their short durations, typical for short GRBs, and their high energy releases, typical for long GRBs. Meanwhile, the GRB rate inferred from high-redshift GRBs also remarkably exceeds the prediction of the collapsar model, with an ordinary star formation history. We show that all these contradictions could be eliminated naturally, if we ascribe some high-redshift GRBs to electromagnetic bursts of superconducting cosmic strings. High-redshift GRBs could become a reasonable way to test the superconducting cosmic string model, because the event rate of cosmic string bursts increases rapidly with increasing redshifts, whereas the collapsar rate decreases. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.1278_arXiv.txt": {
        "abstract": "The spherical collapse model is often used to follow the evolution of overdensities into the nonlinear regime. We describe the correct approach to be used in coupled dark energy cosmologies, where a fifth force, different from gravity and mediated by the dark energy scalar field, influences the collapse. We reformulate the spherical collapse description by deriving it directly from the set of nonlinear hydrodynamical Navier-Stokes equations. By comparing with the corresponding relativistic equations, we show how the fifth force should be taken into account within the spherical collapse picture and clarify the problems arising when an inhomogeneous scalar field is considered within a spherical collapse picture. We then apply our method to the case of coupled quintessence, where the fifth force acts among cold dark matter particles, and to growing neutrino quintessence, where the fifth force acts between neutrinos. Furthermore, we review this method within standard cosmologies and apply our analysis to minimally coupled quintessence. We also check past results for early dark energy parametrizations. ",
        "introduction": "A wide variety of theoretical cosmological models can be challenged and discriminated thanks to predictions on structure formation. At the nonlinear level the behavior of $\\Lambda$CDM cosmologies, in which the role of dark energy (DE) is played by a cosmological constant, can significantly differ from dynamical dark energy models. In more realistic scenarios, these allow for DE couplings to other species. Interacting dark energy cosmologies include: coupled quintessence (DE evolution is coupled to dark matter) \\cite{amendola_2000, mangano_etal_2003, amendola_quercellini_2003, amendola_2004, wang_etal_2007, pettorino_baccigalupi_2008, quartin_etal_2008, boehmer_etal_2008, bean_etal_2008, lavacca_etal_2009}; growing neutrino cosmologies \\cite{amendola_etal_2007, wetterich_2007, mota_etal_2008} and MaVaNs (\\cite{fardon_etal_2004, afshordi_etal_2005, bjaelde_etal_2007, brookfield_etal_2006, brookfield_etal_2007, takahashi_tanimoto_2006} and references therein) (DE is interacting with neutrinos); so-called modified gravity theories such as scalar-tensor theories, including F(R) and extended quintessence \\cite{hwang_1990, hwang_1990_2, wetterich_1995, uzan_1999, faraoni_2000, riazuelo_uzan_2002, perrotta_etal_2000, boisseau_etal_2000, perrotta_baccigalupi_2002, pettorino_etal_2005}. In all these cosmologies a fifth force is present, acting on species whose evolution is coupled to the DE evolution. The presence of a fifth force, mediated by the DE scalar field (the cosmon, seen as the mediator of a cosmological interaction) can modify structure formation in a significant way \\cite{baldi_etal_2009, maccio_etal_2003}, in particular at large scales \\cite{wintergerst_etal_2009}. In view of future data, it is therefore important to understand how these theories behave when density perturbations reach nonlinearity.  While up to now N-body simulations represent the best way to numerically evolve structures, other semi-analytical methods have been used to follow perturbations into the nonlinear regime, either using spherical collapse \\cite{peebles_1967, gunn_gott_1972, padhmanabhan, peacock, bilic_etal_2003, pace_etal_2010} or other alternative methods \\cite{peacock_dodds_1996, pietroni_2008, angrick_bartelmann_2010}. In particular, spherical collapse has been used in several occasions in literature for $\\Lambda$CDM \\cite{padhmanabhan, peacock, engineer_etal_1998}, minimally coupled quintessence models \\cite{wang_steinhardt_1998, mainini_etal_2003, mota_vandebruck_2004, maor_lahav_2005, wang_2005, dutta_maor_2006, mota_etal_2007, abramo_etal_2007, creminelli_etal_2009}, coupled quintessence \\cite{mainini_bonometto_2006, nunes_mota_2006} and when parametrizing early dark energy contributions \\cite{bartelmann_etal_2006, sadeh_etal_2007, francis_etal_2008}.  In this paper we give a detailed description of the spherical collapse method and clarify some tricky issues in its applications. We lay particular focus on the calculation of the extrapolated linear density contrast at collapse $\\delta_c$, a quantity of major interest within a spherical collapse description, often used in a Press-Schechter \\cite{press_schechter_1974} approach to estimate dark matter halo mass distributions.  After reviewing results for standard cosmologies like $\\Lambda$CDM, we consider the case in which a fifth force is present in addition to standard gravitational attraction, as in the case of all the interacting dark energy models mentioned above. The inclusion of the fifth force within the spherical collapse picture requires particular attention. Spherical collapse is intrinsically based on gravitational attraction only and cannot account for other external forces unless it is suitably modified. The dynamics in the spherical collapse model are governed by Friedmann equations. Hence, only gravitational forces determine the evolution of the different scale factors and, in turn, of the density contrast.  A detailed comparison between the linearized spherical collapse picture and the linear relativistic equations allows us to first identify the presence (or absence), in the spherical collapse picture, of terms which are a direct signature of the coupling already at the linear level. We show how spherical collapse necessitates to be suitably modified whenever an additional force other than gravity is present and is big enough to influence the collapse. We use this comparison also to show that a standard treatment of spherical collapse may lead to problems even in the uncoupled case when treating inhomogeneities in the scalar field.  A modification of the spherical collapse picture is indeed possible, via a nonlinear analysis of the model. We derive the set of second order differential equations for the density contrast from the nonlinear Navier-Stokes equations described in \\cite{wintergerst_etal_2009}, extending an idea from \\cite{wintergerst_2009}. We show how $\\delta_c$ can be evaluated directly from these equations and how they can serve as a starting point for a reformulation of spherical collapse. Our results match the numerical solution of the nonlinear hydrodynamical equations solved as described in \\cite{wintergerst_etal_2009, wintergerst_2009}.  We apply our method to coupled quintessence scenarios where a coupling is present among dark matter particles, comparing our results with alternative methods presented in the past \\cite{mainini_bonometto_2006, nunes_mota_2006}. As a further application of our method, we consider for the first time spherical collapse within growing neutrino models, where an interaction is active among neutrinos: in this case we obtain an extrapolated linear density at collapse $\\delta_c$ which shows an oscillating behavior, a characteristic feature of the interaction.  Finally we confirm results found in \\cite{francis_etal_2008} on spherical collapse and early dark energy (EDE).  In Sec.~\\ref{sec:spher_coll} we recall the spherical collapse model and its applications to standard cosmologies (Sec.~\\ref{standard_cosm}). In Sec.~\\ref{sec:coupled_quint} we focus on spherical collapse in presence of a fifth force, taking the case of coupled quintessence as an example of fifth force cosmologies. In Sec.~\\ref{sec:spher_coll_coupled_inc}, we demonstrate that the standard spherical collapse leads to wrong results when applied to coupled quintessence: indeed, by comparing with full relativistic equations (Sec.\\ref{relativistic_eqts}) we demontrate that the spherical collapse equations lack terms that are essential in the presence of a fifth force (Sec.\\ref{lack_fifthforce}) and can lead to incorrect results when an inhomogeneous DE scalar field is included within this framework (Sec. \\ref{inhomogeneity}). Consequently, in Sec.~\\ref{sec:spher_coll_coupled_cor}, we illustrate how the spherical collapse can be correctly reformulated in coupled scenarios by basing it on the full nonlinear Navier-Stokes equations for the respective model. We further comment on the careful choice of initial conditions in Sec.~\\ref{sec:ini_conds}. We apply the derived formalism to give results for coupled quintessence (Sec.~\\ref{sec:coupled_results}) and growing neutrinos (Sec.~\\ref{sec:nu_results}). Finally, we use the described framework to confirm results found in \\cite{francis_etal_2008} for uncoupled early dark energy. ",
        "conclusions": "Spherical collapse is a semi analytical method often used to estimate the nonlinear evolution of structures without reverting to complex numerical methods like N-body simulations.  After reviewing its application to standard cosmologies, we have considered the case of coupled dark energy cosmologies, in which a fifth force other than gravity modifies the collapse.  We have shown that the inclusion of the fifth force within the spherical collapse picture deserves particular caution. As spherical collapse is intrinsically based on gravitational attraction via the Friedmann equations, it does not account for other external forces unless it is suitably modified.  We have presented a detailed comparison between the linearized standard spherical collapse picture and the linear relativistic equations, whose results are summarized in Table \\ref{tab:sc_rel_comp}. Applying standard spherical collapse equations to coupled dark energy by adding a coupling in the conservation equations is insufficient to describe the coupled dark energy scenario, as the fifth force is still missing entirely from the evolution of the density contrast $\\delta$. Results in Table \\ref{tab:sc_rel_comp} also show that a standard treatment of spherical collapse may lead to problems even in the uncoupled case, whenever the scalar field is treated as inhomogeneous.  We have illustrated in detail how a modification of the spherical collapse picture which correctly accounts for the presence of a fifth force is still possible. We have derived the set of second order differential equations for the density contrast from the fully nonlinear Navier-Stokes equations. We have then shown how $\\delta_c$ can be evaluate directly from these equations and how the spherical collapse formalism can be reformulated starting from them. Most importantly, we have further checked that our results match the numerical resolution of the nonlinear hydrodynamical equations performed as described in \\cite{wintergerst_etal_2009, wintergerst_2009}.  We have applied our procedure to coupled quintessence scenarios, evaluating the extrapolated linear density at collapse for this class of cosmologies and showing how it depends on the coupling $\\beta$. Furthermore, we have for the first time applied the spherical collapse to the case of growing neutrino quintessence, where neutrinos feel a fifth force interaction that can lead them to cluster at very large scales. In this case, we demonstrate how the extrapolated linear density at collapse shows a characteristic oscillating behavior, different from standard dark energy models. In future work, this result could be used within a Press-Schechter \\cite{press_schechter_1974} formalism to estimate neutrino halo mass distributions. We have further commented on the choice of initial conditions, whose choice has to be made with careful attention when dealing with the extrapolated linear density at collapse. Finally we have used our approach to verify results found in \\cite{francis_etal_2008} on spherical collapse and early dark energy (EDE)."
    },
    "1005/1005.1752_arXiv.txt": {
        "abstract": "We propose a symmetrized form of the softened gravitational potential which is a natural extension of the Plummer potential. The gravitational potential at the position of particle $i$ $(x_i,y_i,z_i)$, induced by particle $j$ at $(x_j,y_j,z_j)$, is given by: \\begin{equation*} \\phi_{ij} = -\\frac{G m_j}{|{r_{ij}}^2+ {\\epsilon_{i}}^2+{\\epsilon_{j}}^2|^{1/2}}, \\label{eq:abs} \\end{equation*} where $G$ is the gravitational constant, $m_j$ is the mass of particle $j$, $r_{ij} = | {(x_i - x_j)^2 + (y_i - y_j)^2 + (z_i - z_j)^2} |^{1/2}$ and $\\epsilon_i$ and $\\epsilon_j$ are the gravitational softening lengths of particles $i$ and $j$, respectively.  This form is formally an extension of the Newtonian potential to five dimensions. The derivative of this equation in the $x,y$, and $z$ directions correspond to the gravitational accelerations in these directions and these are always symmetric between two particles.  When one applies this potential to a group of particles with different softening lengths, as is the case with a tree code, an averaged gravitational softening length for the group can be used.  We find that the most suitable averaged softening length for a group of particles is $\\langle {\\epsilon_j}^2 \\rangle = \\sum_j^N m_j {\\epsilon_j}^2 / M$, where $M = \\sum_j^N m_j$ and $N$ are the mass and number of all particles in the group, respectively.  The leading error related to the softening length is $\\mathcal O \\left (\\sum_j {{\\boldsymbol r}_j d(\\epsilon_j}^2)/{{r_{ij}}^{3}} \\right )$, where ${\\boldsymbol r}_j$ is the distance between particle $j$ and the center of mass of the group and ${d(\\epsilon_j}^2) = {\\epsilon_j}^2 - \\langle {\\epsilon_j}^2 \\rangle$. Using this averaged gravitational softening length with the tree method, one can use a single tree to evaluate the gravitational forces for a system of particles with a wide variety of gravitational softening lengths.  Consequently, this will reduce the calculation cost of the gravitational force for such a system with different softenings without the need for complicated forms of softening. We present the result of simple numerical tests. We found that our modification of the Plummer potential works well. ",
        "introduction": "\\label{sec:introduction}  \\cite{Aarseth1963} introduced a Plummer potential \\citep{Plummer1911} for gravitational $N$-body simulations in order to eliminate the divergence of the potential that occurs when two particles undergo a close encounter.  With a Plummer softening, the modified gravitational potential between two particles $i$ and $j$ is given by \\begin{equation} {\\phi}_{ij} = -\\frac{G m_j}{|{{\\boldsymbol r}_{ij}}^2 + {\\epsilon}^2|^{1/2}}, \\label{eq:plummer_pot} \\end{equation} where $G$, ${\\boldsymbol r}_{ij}$, $m_j$, and $\\epsilon$ are the gravitational constant, the distance between particles $i$ and $j$, the mass of particle $j$ and the gravitational softening length, respectively.  The increase in the calculation cost due to this modification of the potential is quite small. Thus, it is widely used in simulations of self-gravitating systems and the GRAPE (GRAvity PipE) series also adopts this form of the gravitational potential \\citep{Sugimoto+1990, Ito+1991, Okumura+1993, Makino+1997, Kawai+2000, Makino+2003, KawaiFukushige2006, Makino2005, Makino2007}.  Due to the wide adoption of the Plummer potential, there have been several studies on the choice of the optimal gravitational softening length \\citep{Merritt1996, Athanassoula+2000, Dehnen2001}.  In simulations of self-gravitating systems, it is necessary to deal with a wide range in both space and time. For example, one needs to cover the range of $\\sim 100~{\\rm Mpc}$ (for the tidal field) to $\\sim 10~{\\rm pc}$ or less (for the structure of the interstellar medium) in galaxy formation simulations . Usually, the outer region of the target galaxy is represented by massive particles in order to reduce the calculation cost \\citep[e.g.,][]{NavarroBenz1991, KatzWhite1993}.  In simulations of the coevolution of the galactic center and star cluster, it is necessary to use different masses for different components, since required accuracies and spatial resolutions are different \\citep{Fujii+2007, Fujii+2010}.  Thus we need to use different gravitational softening lengths for different types of particles.   If one uses the direct summation method to calculate the gravitational force, it is possible to use an arbitrary symmetric form for the softening length; for example, $[({\\epsilon_i}^2+{\\epsilon_j}^2)/2]^{1/2}$, $(\\epsilon_i+\\epsilon_j)/2$, or $\\max(\\epsilon_i,\\epsilon_j)$.  On the other hand, when one uses the tree method \\citep{Appel1985,BarnesHut1986}, one needs to use separate trees for each of the different groups of particles with the same gravitational softening length \\citep[e.g.][]{Springel+2001}, since otherwise there will be an error in the force calculation of the order $\\epsilon^2$. If one were to calculate the force from the center of mass of two particles with widely different gravitational softening values, it is not clear how one would take into account the gravitational softening length. However, if one uses different trees for particles with different gravitational softening lengths, the calculation cost increases by a large factor.   \\cite{HernquistKatz1989} proposed a different approach for softening.  They assumed that each particle has a finite size and mass distribution given by a function with a compact support, e.g. a spline function.  The advantage of this gravity kernel is that it becomes a purely Keplarian potential for distances larger than the particle size.  Thus, for particles separated by distances larger than their sizes, one can use the center-of-mass approximation or the multipole expansion of a purely Keplarian potential.  However, spline functions require a few conditional branches and the evaluation of a polynomial.  These operations lead to a non-negligible slowdown in the force and potential calculations on modern CPUs/GPUs which have SIMD execution units.  As a result, the calculation cost of a spline function is higher than that of the simple Plummer softening.  Moreover, the tree approximation cannot be used for particles within the softening length. Thus, if there are many particles within the softening length, they can cause a large increase in the calculation cost.   In this paper, we propose a symmetrized gravitational potential between two particles with different softenings, which is a natural extension of the Plummer potential.  We derive the multipole expansion of a group of such particles with this symmetrized potential. We then show that, in terms of accuracy, it is the best to use the second moment of the softening lengths of the particles in order to obtain the averaged softening length for the group.  The multipole moment up to the second order has a simple form, which is easy and inexpensive to calculate.   This paper is structured as follows: in section \\ref{sec:sym}, we present a symmetrized gravitational potential for a pair of particles with different gravitational softening lengths and then introduce an averaged gravitational softening length for a group of such particles.  The use of this potential with the tree method is also discussed in the section.  Results of numerical tests are given in section \\ref{sec:tests}.  A summary and discussion are shown in section \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  In this paper, we introduced a natural symmetrization of the Plummer potential. This symmetrization is quite simple and hence, easy for anyone to implement into their own code, providing their code adopts the traditional Plummer potential for gravity.   By applying this potential to a group of particles, we derived an averaged gravitational softening length for the group.  Consequently, we can use this potential in the tree method.  The multipole expansion of the symmetrized Plummer potential allows us to use a single tree to calculate forces and potentials for a system of particles that have different softening lengths. Thus the calculation cost of our method is less than that of previous implementations.   Since the modification is quite simple, the modified Plummer potential can be easily implemented in modern GRAPEs ({\\it i.e.}, GRAPE-7 {\\footnote {GRAPE-7 employed field programmable gate array (FPGA) and constructs the gravity pipelines on runtime \\citep{KawaiFukushige2006}.}} and GRAPE-DR {\\footnote {GRAPE-DR uses the original programmable SIMD processors \\citep{Makino2005, Makino+2007}.}}) as they can modify the force law via software libraries. However, for older GRAPEs, the gravity pipelines are implemented on the hardware layer and so it is impossible to use this modified potential.  The latest version of the GRAPE-DR library supports this new symmetrized Plummer potential.    Pure software libraries, {\\it i.e.} those that run on typical CPUs without special hardware, can easily include our calculation model.  Phantom-GRAPE (Nitadori et al, in preparation) is an optimized library that utilizes the SIMD instructions of X86 CPUs and has application interfaces comparable with GRAPE-5 \\citep{Kawai+2000}. For this software, we can easily modify the gravity kernel in the library. We used this modified Phantom-GRAPE for the tests described in Section \\ref{sec:tests}.   The implementation of our modified Plummer potential on GPGPUs \\citep{Nakasato2009, Hamada+2009} is also straightforward."
    },
    "1005/1005.3847.txt": {
        "abstract": "The Bullet Cluster is a massive galaxy cluster at $z=0.297$ undergoing a major supersonic (Mach 3) merger event.  Using { \\it Spitzer} 24\\micron\\ images, IRAC data, optical imaging, and optical spectroscopy, we present the global star formation rate (SFR) of this unique cluster.  Using a 90\\% spectroscopically complete sample of 37 non-AGN MIPS confirmed cluster members out to $R<1.7$ Mpc, and the \\citet{rieke2009} relation to convert from 24\\micron\\ flux to SFR, we calculate an integrated obscured SFR of 267 \\Msolaryr\\ and a specific star formation rate of 28 \\Msolaryr\\ per $10^{14}$\\Msun.   The cluster mass normalized integrated SFR of the Bullet Cluster is among the highest in a sample of eight other clusters and cluster mergers from the literature.  Five LIRGs and one ULIRG contribute 30\\% and 40\\% of the total SFR of the cluster, respectively.   To investigate the origin of the elevated specific SFR, we compare the infrared luminosity function (IR LF) of the Bullet Cluster to those of Coma (evolved to z=0.297) and CL1358+62.  The Bullet Cluster IR LF exhibits an excess of sources compared to the IR LFs of the other massive clusters.  A Schechter function fit of the Bullet Cluster IR LF yields $L^{*}=44.68\\pm0.11$ ergs s$^{-1}$, which is $\\sim$0.25 and 0.35 dex brighter than $L^{*}$ of evolved Coma and CL1358+62, respectively.  The elevated IR LF of the Bullet Cluster relative to other clusters can be explained if we attribute the ``excess'' star-forming IR galaxies to a population associated with the infalling group that have not yet been transformed into quiescent galaxies.   In this case, the timescale required for quenching star formation in the cluster environment must be longer than the timescale since the group's accretion -- a few hundred million years.  We suggest that ``strangulation'' is likely to be an important process in the evolution of star formation in clusters. ",
        "introduction": "The fraction of star-forming galaxies in galaxy clusters is known to be suppressed relative to what is typically found in the less dense environments of galaxy groups and the general field population.  The overwhelming observational evidence \\citep{dressler1980,gomez2003} suggests that star formation must be quenched at some point during hierarchical mass assembly.  Galaxy transformations are thought to occur predominantly in one of two broad categories -- prior to cluster assembly within the galaxy group environment, also known as galaxy pre-processing \\citep{zabludoff1998,kodama2001}, or within the cluster environment itself.  While there is evidence in favor of both scenarios \\cite[e.g.][]{fujita2004,cortese2006,koyama2008,berrier2009}, in the end the role of the cluster environment on the star formation history of its galaxies is unclear.  Even less clear is the impact of a major cluster merger on galaxy evolution.   \\citet{owen2005,miller2003,ferrari2005}  found evidence for triggered star formation in cluster mergers based on an enhanced fraction of star-forming radio galaxies and a preferential distribution of emission line galaxies between merging subclumps.  In contrast, \\citet{poggianti2004} found that the post-starburst population of dwarf galaxies in the Coma cluster lie near the edges of two merging substructures, suggesting that the merger quenches star formation, though these could be relics of starbursts induced during an earlier phase of the merger \\citep{mahajan2010}.   There has also been much effort placed in understanding the specific physical mechanisms responsible for transforming galaxy properties such as morphology and star formation.  Although there is some observational and/or theoretical support for various physical processes considered to occur in the cluster environment, such as ram pressure \\citep{gunn1972}, strangulation \\citep{kawata2008}, and galaxy harassment \\citep{moore1996}, it is still unknown which of these play a dominant role in transforming a star-forming galaxy into a quiescent one.   This is the second in a series of papers to examine the star formation properties of 1E0657-56, also known as the Bullet Cluster. The Bullet Cluster is a galaxy cluster at $z=0.297$ undergoing a major merger event, with a collision between the main cluster and subcluster occuring close to the plane of the sky with $i<8^{\\circ}$ \\citep{markevitch2004}.  A well-defined bow shock front has been confirmed by \\citet{markevitch2002}, and is propagating through the X-ray gas in the subcluster region at a velocity of 4740$\\pm630$ \\kms~(Mach number $M=3.0\\pm0.4$) \\citep{markevitch2006}.  The subcluster itself lags behind the shock front due to a wind from the main cluster gas and travels at 2700 \\kms\\ relative to the main cluster \\citep{springel2007}.  At this velocity, the time since core passage is $\\sim$0.25 Gyr ago.   With our extensive multi-wavelength dataset and optical spectroscopy, we study star formation out to nearly the virial radius of the cluster {\\it ($R_{200}=2.2$ Mpc)}.  In this series of papers, we will examine various aspects of star formation in the Bullet Cluster to understand how a recent major merger event affects the star formation of cluster galaxies.  We also consider the importance of different physical mechanisms that may dominate galaxy evolution in the diverse environments present in the Bullet Cluster.  Star formation in the core of the Bullet Cluster has been studied by \\citet{chung2009} who have found that the strong ram pressure exerted by supersonic gas from the cluster merger does not have a significant impact on recent star formation in the cluster galaxies.   In this paper we use the {\\it Spitzer} Multiband Imaging Photometer (MIPS) in the 24\\micron\\ band \\citep{rieke2004}, to compare the global star formation rate of the Bullet Cluster, as traced by the dusty star-forming galaxies, to those of other merging and non-merging clusters from the literature.  We also present the infrared luminosity function of the Bullet Cluster as a means of investigating the nature of the infrared luminous galaxy population in this massive major cluster merger. ",
        "conclusions": "We use 24\\micron\\ MIPS data to calculate the mass normalized global star formation rate of 37 spectroscopically confirmed Bullet Cluster members within $R<1.7$ Mpc, excluding AGN sources using IRAC, X-ray, and optical spectroscopic data.  The integrated obscured SFR of the Bullet Cluster is 267 \\Msolaryr\\ with a 0.2 dex uncertainty, using the $L_{24}$-SFR relation from \\citet{rieke2009}.  Normalizing by cluster mass we obtain a specific star formation rate of 28 \\Msolaryr per $10^{14}$\\Msun.  Using the \\citet{dale2002} models with the \\citet{kennicutt1998} SFR calibration, and applying a lower limit on 24\\micron\\ flux comparable to previous works, we obtain a mass normalized integrated SFR of 47 \\Msolaryr per $10^{14}$\\Msun.  We find that the specific SFR of the Bullet Cluster is one of the highest among other merging and non-merging clusters from the literature, with the exception of CL0024+16.  In the Bullet Cluster, a small population of LIRGs and ULIRG contribute 30\\% and 40\\% of the total obscured star formation rate.  These galaxies comprise the bright end of the infrared luminosity function, which are not observed in the IR LFs of evolved Coma or CL1358+62.   A Schechter fit to the IR LF of the Bullet Cluster yields $L^{*}=44.70\\pm0.16$, which is $\\sim$0.25 dex brighter than $L^{*}$ of evolved Coma and 0.35 dex brighter than CL1358+62, a massive cluster at $z=0.3$.  We observe an excess of IR bright sources in the Bullet Cluster along the entire range of infrared luminosities in the IR LF, compared to the IR LFs of evolved Coma and CL1358+62.  One way to explain this is to attribute the excess of IR sources to the galaxy population associated with the group that has recently been accreted into the main cluster $\\sim$250 Myr ago (core passage).   The combined IR LF of evolved Coma and mass-scaled SG1120 is consistent within a factor of $\\sim$1.5 of the observed Bullet Cluster IR LF.  This supports the idea that the galaxies originating from the group population have not yet been processed to match typical cluster levels of star formation rate.  In this case, a transformation mechanism such as strangulation, which acts on timescales longer than the merger timescale of the main cluster and subcluster ($\\gap250$ Myr) is necessary to eventually truncate the bright end and suppress the faint end of the IR LF to match the observed distribution in Coma and CL1358+62.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%% TABLES %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{deluxetable}{cccc} \\tabletypesize{\\scriptsize} \\tablecaption{Data} \\tablewidth{0pt} \\tablehead{ \\colhead{Name} & \\colhead{MIPS} & \\colhead{MIPS AGN} & \\colhead{ALL} } \\startdata Confirmed members  &  44  &  6  &  406 \\\\ Interlopers & 83   &  5 &  495 \\\\ Candidate members & 14 & 2 &  ...\\\\ \\enddata \\label{table:data} \\tablecomments{The number of galaxies that are spectroscopicaly confirmed cluster members, interlopers and cluster candidates (based on optical and IRAC colors). The columns show number of galaxies from the MIPS catalog with WFI and IRAC counterparts, number of AGN within the MIPS catalog determined by IRAC colors, X-ray emission, and optical spectroscopy, and the total number of galaxies from the entire IMACS spectroscopic catalog.}  \\end{deluxetable} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{deluxetable*}{ccccc} \\tabletypesize{\\scriptsize} \\tablecaption{LIRGs and AGN} \\tablewidth{0pt} \\tablehead{ \\colhead{Object Type} & \\colhead{RA [hh:mm:ss]} & \\colhead{Dec [$^{\\circ}\\ \\arcmin\\ \\arcsec$]} & \\colhead{$L_{IR}$ [\\Lsun]} &  \\colhead{SFR [\\Msolaryr] } } \\startdata ULIRG/LINER & 6:58:30.87  &  -56:03:36.31 & $1.3\\times 10^{12}$   &   104  \\\\ LIRG & 6:59:13.45  &  -55:56:31.19 & $2.5\\times 10^{11}$   &   28 \\\\ LIRG & 6:58:55.33  &  -55:55:43.04 & $1.5\\times 10^{11}$   &   16  \\\\ LIRG & 6:58:29.30  &  -55:53:15.92 & $1.4\\times 10^{11}$   &   14  \\\\ LIRG & 6:57:47.86  &  -55:55:10.83 & $1.1\\times 10^{11}$   &   11  \\\\ LIRG &  6:58:37.06  &  -56:00:45.66  &  $1.0\\times 10^{11}$  & 10 \\\\ LIRG, X-ray AGN& 6:58:03.52  &  -56:01:14.23 &  $1.0\\times 10^{11}$  & ... \\\\ X-ray \\& IRAC AGN &  6:58:35.22  &  -56:01:4.78  & $8.7\\times 10^{10}$  & ... \\\\ X-ray AGN &  6:58:17.48  &  -56:02:49.35  &  $8.1\\times 10^{10}$  &  ... \\\\ Optical AGN (Seyfert) &  6:58:42.50  &  -56:00:28.22 & $7.6\\times 10^{10}$ &  ... \\\\ IRAC AGN &  6:58:08.48  &  -55:53:36.45 &  $1.7\\times 10^{10}$  &  ...  \\\\ X-ray AGN & 6:58:26.67  &  -56:00:0.18  &  $9.2\\times 10^{9}$ &  ... \\\\  \\enddata \\label{table:LIRGs} \\tablecomments{Spectroscopically confirmed members with MIPS emission that are either LIRGs and/or AGN.  Column 1 notes whether a galaxy is a LIRG/ULIRG, and if it is detected as either an X-ray point source or an IRAC AGN wedge source.  Columns 2-5 show the right ascension, declination, total infrared luminosity, and star formation rate for each galaxy.  SFR of the ULIRG assumes that 35\\% of the total IR luminosity is powered by AGN, and the rest by star formation.  For the LIRGs and AGN, the 1$\\sigma$ error of the measured 24\\micron\\ flux  is between 1-10\\%, and thus the final error of the total infrared luminosity is dominated by the uncertainty in the $L_{24}$ to $L_{IR}$ calibration, which is $\\lap$0.15 dex.} \\end{deluxetable*}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{deluxetable}{cccc} \\tabletypesize{\\scriptsize} \\tablecaption{Schechter Parameters} \\tablewidth{0pt} \\tablehead{ \\colhead{Cluster/group} & \\colhead{$\\alpha$} & \\colhead{$\\log L^{*}$} & \\colhead{$\\phi^{*}$ [N/Mpc$^{2}$/log $L_{IR}$]} } \\startdata Coma (evolved $z=0.3)$   &  -1.4  &  $44.44^{+0.27}_{-0.24}$  &  ... \\\\ CL1358+62 & -1.4  & $44.33^{+0.32}_{-0.25}$ &  $1.92^{+1.4}_{-1.3}$ \\\\ SG1120 & -1.4 & $44.99^{+0.19}_{-0.19}$ &  $1.25^{+0.5}_{-0.3}$   \\\\ Bullet Cluster & -1.4 & $44.68^{+0.11}_{-0.11}$ & $1.77^{+0.5}_{-0.5}$ \\\\ Bullet Cluster (with AGN) & -1.4 & $44.81^{+0.13}_{-0.13}$ &  $1.36^{+0.4}_{-0.4}$ \\\\ \\enddata \\label{table:IRLF} \\tablecomments{Schechter parameters for the clusters and SG1120 shown in Figure~\\ref{fig:IRLFv1}.  Schechter parameters for Coma are adopted from \\citet{bai2009} and evolved to $z=0.3$ using $L^{*}\\propto(1+z)^{3.2}$ evolution, while values for CL1358+62 and SG1120 are from \\citet{tran2009}.  We fit a Schechter function to the observed Bullet Cluster IR LF, without and with the AGN, fixing the faint end slope $\\alpha$ to the match that of Coma.} \\end{deluxetable}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%% FIGURES %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{figure} \\epsscale{1} \\plotone{fig1.eps} \\caption{Spectroscopic redshift distribution with shading to indicate cluster member region.  Inset shows zoomed-in view of cluster members as defined by the caustic analysis of the infall region.} \\label{fig:veldisp} \\end{figure}   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{figure*} \\epsscale{1} \\plottwo{fig2a.eps}{fig2b.eps} \\caption{The R-[4.5] and B-R colors of spectroscopically confirmed cluster members (dots), cluster members with MIPS emission (solid circle), cluster candidates with MIPS emission (open circle), and known interlopers with MIPS emission (cross).   The spectroscopically confirmed cluster galaxies, particularly those with MIPS emission, form a tight sequence in this color-color space.  To minimize contamination from background MIPS sources, we include galaxies only within the diagonal boundaries indicated.  Galaxies with $B-R>3$ are also excluded since these are optically redder than the red sequence.  The right panel shows the R-[4.5] and B-R colors of confirmed cluster members only.  Five confirmed cluster LIRGs are shown as blue stars, and one ULIRG as a green star with a square border.} \\label{fig:select} \\end{figure*}   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{figure*} \\epsscale{1} \\plottwo{fig3a.eps}{fig3b.eps} \\caption{IRAC color-color diagram, showing the AGN wedge \\citep{lacy2004,stern2005}.  Symbols are the same as for Figure~\\ref{fig:select}.  MIPS sources identified as IRAC or X-ray AGN are shown with a blue square and plus sign overplotted on a solid or open red circle.  The Seyfert galaxy identified in Figure~\\ref{fig:BPT} is shown with a blue diamond behind the ULIRG (green star symbol). The dotted line shows the M82 track starting from $z=0$ then progressing to $z=0.3, 0.5, 1,$ and 2, with each redshift marked by a cross.  It is clear that many of the high redshift background sources are removed from our color selection based on Figure~\\ref{fig:select}.} \\label{fig:irac_colors} \\end{figure*}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %BPT FIGURES  \\begin{figure*} \\epsscale{1} \\plottwo{fig4a.eps}{fig4b.eps} \\caption{BPT diagrams with dotted and dashed lines separating purely star-forming galaxies from Seyferts and LINERs \\citep{kewley2006}. Small gray circles are objects that are beyond the IRAC FOV, and therefore it is unknown whether they would fall in the IRAC AGN wedge.  Large circles represent objects within the IRAC FOV, with red (gray) and black circles indicating galaxies with and without 24\\micron\\ emission, respectively.  Objects outlined with a large square and a large triangle have been identified as AGN via IRAC colors and X-ray emission, respectively.  The star symbol represents the ULIRG, which falls close to the Seyfert/LINER boundary in the left panel, and is classified as a LINER in the right panel.} \\label{fig:BPT} \\end{figure*}   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure*} \\epsscale{1.} \\plottwo{fig5a.eps}{fig5b.eps} \\caption{The left and right panels show a comparison of the total infrared luminosity and star formation rate derived from the \\citet{rieke2009} relation (red triangles) compared to those from the \\citet{dale2002} templates with the \\citet{kennicutt1998} star formation relation (black circles), as a function of 24\\micron\\ flux for the 49 non-AGN MIPs confirmed members and cluster candidates in our sample.  The bottom portion of both panels show the systematic offsets between the two methods.} \\label{fig:dale_rieke} \\end{figure*}   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{figure} \\epsscale{1} \\plotone{fig6.eps} \\caption{The star formation rate per stellar mass (specific star formation rate) as a function of stellar mass, with the 50\\% SFR completeness limit indicated (dotted line).  The five LIRGs and one ULIRG are highlighted as star symbols, with a square border around the ULIRG.} \\label{fig:SSFR} \\end{figure}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{figure} \\epsscale{1} \\plotone{fig7.eps} \\caption{The ULIRG imaged with HST/ACS in the F606W filter.  It is a barred spiral with no signs of recent major interactions.} \\label{fig:HST} \\end{figure}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{figure} \\epsscale{1} \\plotone{fig8.eps} \\caption{The optical color-magnitude diagram with spectroscopically confirmed cluster members (dot), spectroscopically confirmed members with MIPS emission (solid circle) and cluster candidates with MIPS emission (open circle).  MIPS sources identified as AGN with IRAC, X-ray, or optical data are indicated with a blue square, triangle, and diamond with a plus sign, respectively. Five spectroscopically confirmed LIRGs are shown as blue stars, and one ULIRG/LINER as a green star with a square border.  Grey solid line shows the color-magnitude relation with slope and normalization adopted from \\citet{lopez-cruz2004}.} \\label{fig:CMD} \\end{figure}   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{figure} \\epsscale{1} \\plotone{fig9.eps} \\caption{The mass normalized integrated SFR as a function of redshift, with the Bullet Cluster shown as a blue solid circle.   SFRs calculated by \\citet{haines2009} for A1758 is illustrated with a triangle. All others are calculated in \\citet{geach2006} and references therein.} \\label{fig:SFR} \\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{figure*} \\epsscale{1} \\plottwo{fig10a.eps}{fig10b.eps} \\caption{The infrared luminosity function of the Bullet Cluster (solid circle), Coma cluster (solid square), CL1358+62 (solid triangle), and SG1120 (asterisk), with the IR LF of the Coma cluster evolved to $z\\sim0.3$.  All infrared luminosities shown here are based on the \\citet{dale2002} galaxy templates. The Schechter fit to the Bullet Cluster IR LF is shown as the dashed curve, and the MIPS 80\\% completeness limit is indicated as a dotted vertical line.  The Left panel shows the Bullet Cluster IR LF excluding all known AGN candidates and the right panel includes six AGN confirmed members and two AGN candidate members (scaled by 0.35 to account for probability of being actual cluster members).  The inclusion of just a small number of AGN can significantly elevate the IR LF, preferentially in the bright end.  However, we note that even without the AGN, the Bullet Cluster IR LF exhibits an excess of IR luminous sources relative to Coma and CL1358+62.  The brightest Bullet Cluster galaxy is a ULIRG and is not included in the Schechter fit.} \\label{fig:IRLFv1} \\end{figure*}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{figure} \\epsscale{1} \\plotone{fig11.eps} \\caption{The infrared luminosity function of the Bullet Cluster (solid circle) excluding known AGN.  Overplotted with solid (red) square is the IR LF of Coma evolved to $z=0.3$ added with the IR LF of SG1120, after the latter has been scaled down by a factor of $\\sim$4 to match the approximate mass of the infalling group population in the Bullet Cluster.} \\label{fig:IRLFv2} \\end{figure}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{figure} \\epsscale{1} \\plotone{fig12.eps} \\caption{R-band WFI image shown with X-ray surface brightness contours.  The large box shows the FOV of our 24\\micron\\ MIPS data, and small boxes indicate the location of the five LIRGs and one ULIRG.} \\label{fig:spatialdist} \\end{figure}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1005/1005.0048_arXiv.txt": {
        "abstract": "{Acoustic peaks in the spectrum of the cosmic microwave background in spherically symmetric inhomogeneous cosmological models are studied. At the photon-baryon decoupling epoch, the universe may be assumed to be dominated by non-relativistic matter% , and thus we may treat radiation as a test field in the universe filled with dust which is described by the Lema\\^itre-Tolman-Bondi (LTB) solution. First, we give an LTB model whose distance-redshift relation agrees with that of the concordance $\\Lambda$CDM model in the whole redshift domain and which is well approximated by the Einstein-de Sitter universe at and before decoupling. We determine the decoupling epoch in this LTB universe by Gamow's criterion and then calculate the positions of acoustic peaks. Thus obtained results are not consistent with the WMAP data. However, we find that one can fit the peak positions by appropriately modifying the LTB model, namely, by allowing the deviation of the distance-redshift relation from that of the concordance $\\Lambda$CDM model at $z>2$ where no observational data are available at present. Thus there is still a possibility of explaining the apparent accelerated expansion of the universe by inhomogeneity without resorting to dark energy if we abandon the Copernican principle. Even if we do not take this extreme attitude, it also suggests that local, isotropic inhomogeneities around us may seriously affect the determination of the density contents of the universe unless the possible existence of such inhomogeneities is properly taken into account. } ",
        "introduction": "\\label{sec:intro}  In the study of cosmology, we usually assume that we do not live in a special position in the universe, the so-called Copernican principle. However, although this principle is natural, it should not be blindly assumed but should be justified observationally. Conventionally observational data are interpreted under the assumption of a homogeneous and isotropic universe on average. Therefore it is not clear at all how big the systematic errors would be in the determination of the cosmological parameters if this assumption were abandoned. In other words, it is important to investigate possible ``anti-Copernican'' models of the universe and test if such models can be observationally excluded. In this paper, as one of such attempts, we try to construct a cosmological model without dark energy which is consistent with the observed distance-redshift relation as well as with the WMAP data.  In anti-Copernican models, we are assumed to be located at a special place in the universe, usually at the center of a spherically symmetric inhomogeneous universe~\\cite{Zehavi:1998gz, Tomita:1999qn,Tomita:2000jj,Tomita:2001gh,Celerier:1999hp}. In recent years, such models have attracted much attention~\\cite{Alnes:2005rw,Alnes:2006pf,Alnes:2006uk, Alexander:2007xx,Bolejko:2008cm,Enqvist:2007vb,Enqvist:2006cg, February:2009pv,GarciaBellido:2008nz,GarciaBellido:2008yq, GarciaBellido:2008gd,Garfinkle:2006sb,Kasai:2007fn, 2000ApJ...529...26T,Tomita:2000rf,Zibin:2008vk,Regis:2010iq,Clifton:2009kx, Kodama:2010gr}, and various ways to observationally test these models have been proposed by many authors~\\cite{Biswas:2007gi, Bolejko:2005fp,Bolejko:2008xh,Bolejko:2008ya, Brouzakis:2006dj,Brouzakis:2007zi,Bhattacharya:2009bz,Clarkson:2007pz, Caldwell:2007yu,Clifton:2008hv,Dabrowski:1997sm,Goodman:1995dt, Godlowski:2004gh,Jia:2008ti,Lasky:2010vn,Marra:2007pm,Moffat:2006ct, PascualSanchez:1999zr,Quartin:2009xr,Romano:2007zz,Stelmach:2006zc, Tanimoto:2009mz,Tomita:2009yx,Uzan:2008qp}.  One of the simplest ways to construct an anti-Copernican model is to solve the inverse problem of the distance-redshift relation. Although the isotropy of the universe around us has been confirmed with high accuracy by the observation of the cosmic microwave background (CMB), this does not automatically imply homogeneity of the universe. Thus, in solving the inverse problem, we may assume that the universe is spherically symmetric around us. In addition, we usually assume that the universe is dominated by cold dark matter, that is, by dust. The spherically symmetric dust filled spacetime is described by the Lema\\^itre-Tolman-Bondi (LTB) solution. The LTB solution has three arbitrary functions of the radial coordinate approximately corresponding to the density profile, the spatial curvature and the big-bang time perturbation, with one of them being a gauge degree of freedom representing the choice of the radial coordinate. These arbitrary functions may be determined by requiring that the resulting LTB universe be consistent with selected important observational data (e.g., the distance-redshift relation). However, we should note that it is not apparent at all if these three functions have enough degrees of freedom to fit all of the important observational data.  In 1999, C\\'{e}l\\'{e}rier solved the inverse problem analytically at small redshifts $z\\ll1$ in the form of the Maclaurin series~\\cite{Celerier:1999hp}. Then, in 2002, Iguchi, Nakamura and Nakao constructed numerically an LTB model whose distance redshift relation agrees with that of the $\\Lambda$CDM model at $z\\lesssim1.6$~\\cite{Iguchi:2001sq} \\footnote{Before Refs.~\\cite{Celerier:1999hp,Iguchi:2001sq}, Mustapha et.al discussed an inverse problem using the LTB solution~\\cite{Mustapha:1998jb}. }. However, they could not go beyond $z\\sim1.6$ due to a technical problem. Since then the inverse problem has been discussed by various authors~\\cite{Chung:2006xh,Vanderveld:2006rb,Celerier:2009sv, Kolb:2009hn,Romano:2009mr,Romano:2009ej,Yoo:2008su}. In 2008, Yoo, Kai and Nakao succeeded in constructing an LTB model whose distance-redshift relation agrees with that of the concordance $\\Lambda$CDM model in the whole redshift domain and which is homogeneous in the early stage of the universe, that is, the big-bang time perturbation being taken to be zero~\\cite{Yoo:2008su}. Note that because of the existence of a gauge degree of freedom in the choice of the radial coordinate, there remains only one functional degree of freedom. In \\cite{Yoo:2008su}, this is represented by the function that determines the spatial curvature.  Not only the distance-redshift relation but also CMB observations in inhomogeneous cosmology have been widely discussed~\\cite{Alnes:2006pf,Alnes:2005rw,Alexander:2007xx, Bolejko:2008cm,Clarkson:1999yj,Caldwell:2007yu,GarciaBellido:2008gd, GarciaBellido:2008nz,GarciaBellido:2008yq,Goodman:1995dt, Godlowski:2004gh,Zibin:2008vj,Zibin:2008vk,Regis:2010iq,Clifton:2009kx,Kodama:2010gr}. In many of these works, it is assumed that the universe is homogeneous in the spacelike asymptotic region from which the CMB photons come, hence the CMB photon distribution is assumed to be the same as that in the homogeneous and isotropic universe models at the time of decoupling. Several authors proposed parametrized LTB models and gave constraints on the parameters with observational data~\\cite{Alnes:2005rw,Alexander:2007xx, Bolejko:2008cm,Enqvist:2007vb,GarciaBellido:2008nz, GarciaBellido:2008yq,GarciaBellido:2008gd,Zibin:2008vk}.   Recently, Bolejko and Wyithe~\\cite{Bolejko:2008cm} suggested that it is possible to construct an LTB model which is consistent not only with the distance-redshift relation but also with the acoustic peak positions of the WMAP data~\\cite{Hinshaw:2006ia} by choosing appropriately the arbitrary functions of the LTB solution. In this paper, by modifying the LTB model given in Ref.~\\cite{Yoo:2008su} we show explicitly that this is indeed possible. For simplicity, we use Gamow's criterion to determine the decoupling epoch instead of invoking a precise numerical calculation. This is the same procedure adopted in Ref.\\cite{Regis:2010iq}. An advantage of this simplified prescription is that it makes it easy to understand the physical degrees of freedom in the LTB universe that determine the CMB anisotropy spectrum. We note that, in our construction, only the asymptotic homogeneity on the past light cone of the observer at the symmetry center was necessary in contrast to a much stronger asymptotic spatial homogeneity condition (see a related discussion in Ref.~\\cite{Enqvist:2009hn}).   This paper is organized as follows. In \\S\\ref{sec:background}, we give a brief review of the LTB solution and give a fitting function for the spatial curvature of the LTB model that has the exactly same distance-redshift relation as the concordance $\\Lambda$CDM model~\\cite{Yoo:2008su}. In \\S\\ref{sec:LSSinLTB}, we discuss Gamow's criterion and the physical degrees of freedom at decoupling. The position of the first acoustic peak in the CMB spectrum in an LTB universe is discussed in \\S\\ref{sec:comparison}. In \\S\\ref{sec:inverseproblem}, we construct an LTB model which is consistent with the acoustic peak positions of the WMAP data by modifying the asymptotic structure of the LTB model obtained in~\\cite{Yoo:2008su}. \\S\\ref{sec:sudi} is devoted to summary and discussion. ",
        "conclusions": "\\label{sec:sudi} In this paper, we investigated if an LTB universe model can account not only for the observed distance-redshift relation but also for the observed peak positions in the CMB spectrum, assuming that we are at the symmetry center of the universe.  First, we presented an LTB model whose distance-redshift relation is equal to that of the concordance $\\Lambda$CDM model in the whole redshift domain. This model is expressed in terms of a fitting function for the curvature numerically obtained in Ref.~\\cite{Yoo:2008su}. Then, we determined the last scattering surface for photons in this model by numerically integrating the past-directed radial null geodesics emanating from the central observer until Gamow's criterion is satisfied.  The metric of this LTB model can be well approximated by the EdS universe at the decoupling epoch. Thus, assuming that our model universe experiences the same history as that of a homogeneous and isotropic universe before decoupling, we evaluated the positions of the acoustic peaks in the CMB spectrum. It was found that thus obtained position of the first acoustic peak deviates from the observed position more than 7\\%, implying that our LTB universe model whose distance-redshift relation agrees with that of the concordance $\\Lambda$CDM model for the entire redshift domain is ruled out.  In order to resolve this inconsistency, we considered modifications of the LTB model in the region far from the symmetry center, corresponding to the redshift domain of $z>2$. Since the observation of Type Ia supernovae are limited within $z<2$, this modification does not contradict the current observational data. Then we considered a one-parameter family of LTB models whose distance-redshift relation matches the observational data at $z<2$, and found a parameter that gives the first acoustic peak position consistent with the observed peak position. Then we checked the second and third peak positions and found that they are also consistent with the WMAP data~\\cite{Hinshaw:2006ia}.  We have not discussed relative heights of acoustic peaks in this paper. At the moment, we have no idea if one can also fit them by an LTB model. There is, however, a hint in the literature. In Ref.~\\cite{Blanchard:2003du}, it was reported that the observed CMB spectrum can be regarded as that in the EdS universe with a low value of the Hubble constant if one tunes the primordial power spectrum of the density perturbation. If this is the case, then we may be able to find an LTB model which can account for the observed relative heights of the acoustic peaks. This issue is left for future work.  In this paper, we have not introduced isocurvature components in the inhomogeneity on the last scattering surface. In a forthcoming paper, we plan to investigate if we may construct an LTB model consistent with the observation of the kinematic Sunyaev-Zel'dovich effect. Our preliminary analysis indicates that isocurvature perturbations between the non-relativistic matter and photons are necessary in order to explain it. However, if we introduce large radial isocurvature perturbations on the decoupling hypersurface, the history of the inhomogeneous universe model before decoupling would differ substantially from the conventional adiabatic perturbation scenario in the homogeneous and isotropic background. In such a case, our analysis of the acoustic peaks in the CMB spectrum may be invalidated. However, since we have considered only a single-parameter family of LTB models, while there still remains a functional degree of freedom, it is premature to make any definite statement at the moment. Apparently much more work seems necessary to exclude or select the LTB universe as a viable alternative model of our universe."
    },
    "1005/1005.2421_arXiv.txt": {
        "abstract": "The intergalactic medium (IGM) is the dominant reservoir of baryons at all cosmic epochs.  In this paper, we investigate the evolution of the IGM from $z=2\\rightarrow 0$ in $(48\\hmpc)^3$, 110-million particle cosmological hydrodynamic simulations using three prescriptions for galactic outflows. We focus on the evolution of IGM physical properties, and how such properties are traced by \\lya\\ absorption as detectable using {\\it Hubble's} Cosmic Origins Spectrograph (COS).  Our results broadly confirm the canonical picture that most \\lya\\ absorbers arise from highly ionized gas tracing filamentary large-scale structure.  Growth of structure causes gas to move from the diffuse photoionized IGM into other cosmic phases, namely stars, cold and hot gas within galaxy halos, and the unbound and shock-heated warm-hot intergalactic medium (WHIM).  By today, baryons are comparably divided between bound phases (35\\% in our favoured outflow model), the diffuse IGM (41\\%), and the WHIM (24\\%).  Here we (re)define the WHIM as gas with overdensities lower than that in halos ($\\rho/\\bar{\\rho}\\la 100$ today) and temperatures $T>10^{5}$~K, to more closely align it with the ``missing baryons\" that are not easily detectable in emission or \\lya\\ absorption.  Strong galactic outflows can have a noticeable impact on the temperature of the IGM, though with our favoured momentum-driven wind scalings they do not.  When we (mildly) tune our assumed photoionizing background to match the observed evolution of the \\lya\\ mean flux decrement, we obtain line count evolution statistics that broadly agree with available (pre-COS) observations.  We predict a column density distribution slope of $f(\\nhi)\\propto \\nhi^{-1.70}$ for our favoured wind model, in agreement with recent observational estimates, and it becomes shallower with redshift.  Winds have a mostly minimal impact, but they do result in a shallower column density slope and more strong lines.  With improved statistics, the frequency of strong lines can be a valuable diagnostic of outflows, and the momentum-driven wind model matches existing data significantly better than the two alternatives we consider.  The relationship between column density and physical density broadens mildly from $z=2\\rightarrow 0$, and evolves as $\\rho\\propto \\nhi^{0.74} 10^{-0.37z}$ for diffuse absorbers, consistent with previous studies.  Linewidth distributions are quite sensitive to spectral resolution; COS should yield significantly broader lines than higher-resolution data.  Thermal contributions to linewidths are typically subdominant, so linewidths only loosely reflect the temperature of the absorbing gas. This will hamper attempts to quantify the WHIM using broad \\lya\\ absorbers, though it may still be possible to do so statistically.  Together, COS data and simulations such as these will provide key insights into the physical conditions of the dominant reservoir of baryons over the majority of cosmic time. ",
        "introduction": "Diffuse intergalactic hydrogen produces numerous weak redshifted \\ion{H}{i} \\lya~(1216\\AA) absorption features along lines of sight to distant bright objects such as quasars, the phenomenon known as the \\lya\\ forest~\\citep{lyn71,sar80}.  Advances in sensitive high-resolution spectroscopy and theoretical understanding of structure formation have led to the currently favoured paradigm for the origin of the \\lya\\ forest~\\citep{cen94,zha95,mir96,her96}: it arises from highly photoionized intergalactic hydrogen~\\citep{gun65} tracing uncollapsed large-scale structure in the Cosmic Web.  The general properties of the \\lya\\ forest are well-described by the Fluctuating Gunn-Peterson Approximation~\\citep[FGPA;][]{wei97a,cro98}, in which a tight relation between density and temperature in cosmologically-expanding photoionized gas~\\citep{hui97} leads to a tight correlation between the \\lya\\ optical depth and the underlying mass density at each location along the line of sight. At $z\\ga 2$, where the \\lya\\ transition falls into the observed-frame optical, this approximation is quite good, and it can be used to infer the mean baryon density \\citep{rau97,wei97b} and the matter fluctuation spectrum, which in turn provides constraints on cosmological parameters \\citep{cro99,cro02,mcd00,mcd06,vie04,vie08}.  In this paper, a successor to \\cite{dav99} and \\cite{dav01}, we use cosmological hydrodynamic simulations to model the IGM and the \\lya\\ forest at lower redshifts.  Here the \\lya\\ transition falls in the ultraviolet (UV), requiring more difficult space-based observations.  Furthermore, the physical description of the \\lya\\ forest departs more strongly from the FGPA, because filamentary structures begin to grow so large as to heat a significant fraction of intergalactic gas to well above photoionization temperatures \\citep{dav99}.  This results in the emergence of the Warm-Hot Intergalactic Medium (WHIM; \\citealt{cen99,dav99,dav01,cen06}), which contains a substantial fraction of all baryons today.  Additionally, the evolving relationship between matter overdensity and \\lya\\ column density is such that at low-$z$, a given overdensity produces a lower column density absorber: roughly speaking, a $10^{14.5}\\cm^{-2}$ line at $z\\approx 2.5$ is physically analogous to a $10^{13}\\cm^{-2}$ line at $z=0$ \\citep{dav99}. This shift results in many fewer \\lya\\ forest absorbers to a given sensitivity limit, i.e. a much thinner forest of lines. All these factors make the low-$z$ \\lya\\ forest more difficult to characterize both observationally and theoretically, and not as useful for constraining cosmology~\\citep{zha05}.  But the intergalactic medium (IGM) at low-$z$ still contains the vast majority of baryons in the Universe, so understanding the low-$z$ \\lya\\ forest and the WHIM remains a central challenge to assembling a complete picture of cosmic baryon evolution.  One of the original Key Projects of {\\it Hubble Space Telescope} was to understand the low-$z$ \\lya\\ forest.  The Faint Object Spectrograph (FOS), despite being unable to fully resolve \\lya\\ absorbers, provided critical insights into the nature of the low-$z$ \\lya\\ forest.  For example, even the earliest observations showed that the number density evolution of \\lya\\ absorbers must slow dramatically between $z>2$ and $z=0$ \\citep{bah91,mor91}.  Initially, it was suggested that low-$z$ absorbers might represent a different population than those at high-$z$, with the high-$z$ systems tracing the Cosmic Web and the low-$z$ systems arising in galaxy halos~\\citep{bah96}.  This idea was supported by the apparently tight connection between column density and galaxy impact parameter~\\citep{che98}.  However, \\citet{dav99} showed that all these observations can be fully accomodated within the structure formation model of the \\lya\\ forest.  The change in number density evolution can be explained as an ionization effect~\\citep{the98}, since the dominant source of metagalactic flux (quasars) diminishes rapidly at $z\\la 2$~\\citep{haa01}, countering the effect of the declining mean density of baryons. The relationship between \\lya\\ equivalent width and galaxy impact parameter arises because the IGM is denser around galaxies owing to matter clustering.  Subsequent higher-resolution quasar spectroscopy using the Space Telescope Imaging Spectrograph (STIS) has generally confirmed these interpretations, and STIS \\lya\\ forest data is well-described by structure formation models~\\citep{dav01b}. Hence the low-$z$ \\lya\\ forest is not fundamentally different than the high-$z$ one, although the relationship of neutral hydrogen absorption to large-scale structure has shifted.  The installation of the Cosmic Origins Spectrograph (COS) on {\\it Hubble} should usher in a new age for the study of the low-$z$ \\lya\\ forest.  The dramatically increased sensitivity over STIS will enable a more robust characterization of the $z\\la 1$ IGM, and it will enable study of \\lya\\ absorbers at overdensities comparable to that detectable at high-$z$.  COS will also enable tomographic probes of the IGM at transverse separations that will provide interesting constraints on the large-scale coherence of \\lya\\ absorbers~\\citep[e.g.][]{cas08}.  All these data will provide theoretical studies with new opportunities to understand the low-$z$ IGM, and new challenges for models to meet.  An advantage of low-$z$ IGM work compared to high-$z$ is that the galaxy population around low-$z$ absorbers can be studied to much greater depth.  The relation between the two provides unique insights into how the early and late stages of gravitational collapse are connected, and how the baryons have decoupled from the dark matter via shock heating and radiative cooling.  The galaxy-absorber connection can be studied through the individual association of galaxies with absorbers or by statistical measures such as correlation functions.  This is the subject of a subsequent paper in this series (Kollmeier et al., in preparation).  Another paper in this series (Oppenheimer et al., in preparation) will deal with metal line absorbers detectable using COS, and their relationship to galaxies and large-scale structure.  In this paper, we employ state-of-the-art cosmological hydrodynamic simulations to study the $z\\la 1$ IGM as traced by the \\lya\\ forest and WHIM.  Where appropriate we include comparisons with existing data and predictions for COS.  Our new simulations have $216\\times$ more particles and $80\\times$ more volume compared to our previous studies in \\citet{dav99} and \\citet{dav01b}, and they include more sophisticated physical processes compared to those works and more recent studies such as \\citet{pas09}.  For example, our simulations now include metal-line cooling and a well-constrained heuristic model for galactic winds that enrich the IGM and regulate star formation in accord with observations.  While we will not discuss chemical enrichment in this work~\\citep[see e.g.][]{opp08}, the energy input from such outflows can non-trivially impact the thermal conditions in the IGM, particularly around galaxies.  \\citet{tor10} investigated the IGM in simulations similar to ours with the further addition of feedback from black holes, showing that it can significantly impact WHIM properties.  In \\S\\ref{sec:sims} we describe our cosmological hydrodynamic simulations and the generation and analysis of artificial spectra. In \\S\\ref{sec:igmphys} we describe the physical conditions of the IGM from $z=2\\rightarrow 0$, including how \\lya\\ absorbers trace large-scale structure and the emergence of WHIM gas.  In \\S\\ref{sec:stats} we present statistics of \\lya\\ absorbers and their evolution. We also present comparisons to STIS data and predictions for COS, and study the impact of differences in spectral resolution and noise. In \\S\\ref{sec:phys} we discuss how \\lya\\ absorbers trace the physical conditions of the low-$z$ IGM, including wide absorbers that trace WHIM gas.  We summarize and discuss the implications of our results in \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We study the low-redshift IGM as traced by \\lya\\ absorbers in $48\\hmpc$, $2\\times 384^3$-particle cosmological hydrodynamic simulations with three prescriptions for galactic outflows.  This work builds on our previous works of a decade ago, namely \\citet{dav99} and \\citet{dav01}, by employing simulations with $216\\times$ more particles and $80\\times$ more volume, using improved numerical methodology, and most significantly, testing the impact of large-scale galactic outflows on the IGM as traced by \\lya\\ absorption.  Our scenario for the nature and evolution of \\lya\\ forest absorbers from $z=2\\rightarrow 0$ remains basically unchanged: most absorbers arise in low-overdensity, highly photoionized gas that traces uncollapsed large-scale filaments.  \\citet{dav99} argued that much of the observed evolution in the \\lya\\ forest can be understood in terms of the evolving relationship between \\ion{H}{i} column density and the overdensity of the absorbing gas; our latest simulations yield a similar slope and evolution of this relation.  Outflows are found to have a small impact on the diffuse IGM, but they do have a significant effect on bound-phase absorbers that arise nearer to galaxies. Our statistical results for \\lya\\ forest absorbers are obtained by applying the AutoVP line-fitter \\citep{dav97} to artificial spectra with a resolution and S/N similar to those expected for typical COS observations.  Our main results are as follows: \\begin{itemize} \\item Baryons at $z\\sim 0$ are divided comparably amongst the diffuse phase, the WHIM, and the bound phases (condensed, hot halo, and stars), The WHIM here is defined as low-density (outside of halos) gas with $T>10^{5}$~K, a notable change from previous definitions that impose a maximum temperature but not a maximum density. In detail, our favoured momentum-driven outflow model has 41\\% in diffuse gas, 35\\% in bound phases, and 24\\% in WHIM at $z=0$, with values for other wind models and redshifts shown in Figure~\\ref{fig:baryphase}. \\item Our favoured momentum-driven wind (vzw) model shows similar evolution in diffuse (unbound) IGM baryonic phases to the no-wind (nw) case.  In contrast, the condensed, hot halo, and stellar fractions are all noticeably different between those two, and between either of these models and the constant wind (cw) model.  Furthermore, cw significantly heats the IGM, resulting in more WHIM and less condensed and diffuse gas. \\item The vzw model at COS resolution shows a column density distribution with a power-law slope of $-1.70$. The slope is slightly steeper for nw and cw, and it is steeper for higher-resolution spectra.  While formally all slopes are within $\\sim 2\\sigma$ of each other, recall that the lines of sight are identical through each simulation, so the differences are fairly robust. \\item The shallower vzw slope arises because it produces more strong ($\\nhi>10^{14}\\cdunits$) lines at $z=0$, which agrees better with observations of $dN/dz$ evolution.  Weaker lines show no differences in $dN/dz$ among wind models, but they do indicate a discrepancy around $z\\sim 1$ that may reflect deviations from the \\citet{haa01} ionizing background evolution. \\item The linewidth ($b$) distribution shows a median around $\\sim 40\\kms$ at COS resolution, and it is insensitive to the wind model.  It is highly sensitive to spectral resolution, even at large $b$ values, and much of the broadening owes to instrumental resolution (note that our $b$-parameters are not corrected for this). \\item Diffuse \\lya\\ absorbers follow a relation between $\\nhi$ and overdensity given by $\\rho/\\bar\\rho \\propto \\nhi^{0.74}10^{-0.37z}$, in general agreement with our previous findings.  This results in a region of a given overdensity giving rise to a lower column density systems at higher redshift.  Departures from this relation become more prominent at low-$z$ owing to an increased amount of shock heating. \\item The bulk of \\lya\\ absorbers have total hydrogen columns of $\\sim 10^{18-19}\\cdunits$, reflecting the trend that the low-$\\nhi$ absorbers are more ionized (neutral fraction $\\sim 10^{-6}$~cm$^{-3}$) and larger (hundreds of kpc), while high-$\\nhi$ absorbers are less ionized (neutral fraction $\\sim 10^{-3}$~cm$^{-3}$) and smaller (few kpc). \\item Linewidths only loosely reflect the temperature of the absorbing gas, because typically linewidths are $1.5-3$ times the thermal broadening value.  Hence, wide lines do not exclusively track WHIM or hot gas, though trends in $\\rho$ and $T$ with both $b$ and $\\nhi$ could be used to assign absorbers to a particular cosmic phase in a statistical sense. \\item Improved $dN/dz$ statistics for weak lines ($\\nhi < 10^{14}\\cdunits$) will improve determinations of the UV background intensity at $z<2$. Improved statistics at higher column densities can provide strong tests of galactic outflow models. \\end{itemize}  The installation of COS on {\\it Hubble} promises a great leap forward in our understanding of low-redshift \\lya\\ absorption and the IGM from which it arises.  Besides testing the basic paradigm of the cosmic web origin for weak \\lya\\ forest absorbers, COS \\lya\\ data will provide constraints on the baryon content in various cosmic phases, ionizing background evolution, and the nature of galactic outflows. Additional constraints will be provided by the distribution of IGM metals that can be probed to much greater precision with COS (the subject of Paper~II in this series) and by the relation between galaxies and absorption systems (covered in Paper~III). The combination of these constraints with overall \\lya\\ absorber statistics will provide powerful diagnostics for the mechanisms of galactic outflows. Comparisons of theory and observations in the coming years will critically test our prevailing paradigms and prescriptions, thereby providing valuable insights into the evolution of the majority of baryons over cosmic time."
    },
    "1005/1005.0338_arXiv.txt": {
        "abstract": "With a state-of-the-art numerical method for solving the integral-differential equation of radiative transfer, we investigate the flux of the Ly$\\alpha$ photon $\\nu_0$ emergent from an optically thick halo containing a central light source. Our focus is on the time-dependent effects of the resonant scattering. We first show that the frequency distribution of photons in the halo are quickly approaching to a locally thermalized state around the resonant frequency, even when the mean intensity of the radiation is highly time-dependent. Since initial conditions are forgotten during the thermalization, some features of the flux, such as the two peak structure of its profile, actually are independent of the intrinsic width and time behavior of the central source, if the emergent photons are mainly from photons in the thermalized state. In this case, the difference $|\\nu_{\\pm}-\\nu_0|$, where $\\nu_{\\pm}$ are the frequencies of the two peaks of the flux, cannot be less than $2$ times of Doppler broadening. We then study the radiative transfer in the case where the light emitted from the central source is a flash. We calculate the light curves of the flux from the halo. It shows that the flux is still a flash. The time duration of the flash for the flux, however, is independent of the original time duration of the light source but depends on the optical depth of the halo. Therefore, the spatial transfer of resonant photons is a diffusion process, even though it is not a purely Brownian diffusion. This property enables an optically thick halo to trap and store thermalized photons around $\\nu_0$ for a long time after the cease of the central source emission. The photons trapped in the halo can yield delayed emission, of which the profile also shows typical two peak structure as that from locally thermalized photons.  Possible applications of these results are addressed. ",
        "introduction": "The transfer of Ly$\\alpha$ radiation in optically thick medium is fundamentally important for the understanding of the physics of halos around Ly$\\alpha$ photon sources or clouds nearby these sources. It includes Ly$\\alpha$ forest, damped Ly$\\alpha$ system, Ly$\\alpha$ blob, Ly$\\alpha$ emitter and fluorescent Ly$\\alpha$ emission of galaxies and quasars, as well as the optical afterglow of gamma ray bursts. The profiles of the emission and absorption of the Ly$\\alpha$ radiation from these sources are powerful tools to constrain the mass density, velocity, temperature and the fraction of neutral hydrogen of IGM at various redshifts (Miralda-Escude \\& Rees 1998; Miralda-Escude 1998; Haiman \\& Cen 2005; Tasitsiomi 2006; Totani et al. 2006; McQuinn et al. 2007).  It is well known that the resonant scattering of Ly$\\alpha$ photons with neutral hydrogen atoms has a profound effect on the time, space and frequency dependencies of the transfer of Ly$\\alpha$ photons. An analytical solution of the integro-differential equation of the resonant radiative transfer revealed that the resonant scattering leads to a local Boltzmann distribution of photons in a small frequency range around the Ly$\\alpha$ frequency $\\nu_0$ (Wouthuysen 1952; Field 1958, 1959). The temperature of the Boltzmann distribution is equal to the kinetic temperature of the neutral hydrogen atoms. The width of the local Boltzmann distribution is increasing with time. This is the so-called Wouthuysen-Field effect, which is important in $21$ cm cosmology (Fang 2009).  Besides Field's solution, no other time-dependent analytical solutions of the integro-differential equation are available. All other analytical solutions (Harrington 1973; Neufeld 1990; Dijkstra et al. 2006) are based on the time-independent Fokker-Planck equation, which is the diffusion approximation of the radiation transfer. These time-independent solutions are important, but can only be used to describe the ``limiting asymptotic behavior\" of the radiative transfer (Adams 1972). It gives nothing of the time-development of the local Boltzmann distribution. The Monte Carlo numerical method has also been used to solve the radiative transfer of Ly$\\alpha$ photons (Lee 1974, Zheng \\& Miralda-Escude 2002; Ahn 2002; Cantalupo et al. 2005; Verhamme et al. 2006). Yet, the time-development of the Wouthuysen-Field effect is also absent in this approach.  Time-independent approximation would be reasonable if the length scale $l$ and time scales $t$ of the problem considered satisfy the condition $l/t \\ll c$. In this case, one can take $c\\rightarrow \\infty$, and then, the time derivative term $\\partial/c\\partial t$ of the radiative transfer equation can be dropped. The condition $l/t \\ll c$, however, cannot be satisfied if the size of the neutral hydrogen halo is large and the time scale of the Ly$\\alpha$ photon source is small. For instance, the time scale $t_{\\rm after}$ of the optical afterglow of GRBs is of the order of a few tens of hours (e.g. Tanvir et al. 2009; Salvaterra et al. 2009), or a few days (Vreeswijk et al. 2004), while the column number density of neutral hydrogen is as large as 10$^{21-22}$ cm$^{-2}$, and number volume density is about 10$^2$ -10$^4$ cm$^{-3}$ (Vreeswijk et al. 2004). Therefore, the size of the neutral hydrogen halos should be much larger than $ct_{\\rm after}$. It is improper to treat the short-time problems with the solutions of ``limiting asymptotic behavior\".  Recently, a state-of-the-art numerical solver for the kinetic equations has been developed.  This solver is based on the weighted essentially non-oscillatory (WENO) scheme (Jiang \\& Shu 1996). It has been developed to solve the Boltzmann equations (Carrillo et al. 2003, 2006) and radiative transfer (Qiu et al. 2006, 2007, 2008). This numerical solver has successfully passed the tests of conservation of photon number, Field's solutions and the Wouthuysen-Field effect etc., and has been properly used as the solver of the transfer of resonant photons (Roy et al. 2009a, b, c).  We will study, in this paper, the time-dependent behavior of the Ly$\\alpha$ radiation transfer in an optically thick medium. We will not try to explain any specific observed Ly$\\alpha$ spectrum, instead we will study the physical features of the resonant photon transfer, which can not be addressed with previous solvers. For instance, we will show that the frequency distribution of the Ly$\\alpha$ photons can keep in a locally thermalized state even when the intensity and flux are highly time-dependent. This feature is essentially important, as thermalization generally will lead to the initial conditions being forgotten. Our solver is also able to study the transfer of a light flash in optically thick halo. It shows that the nature of the transfer is a diffusion process, though it is not a purely Brownian diffusion. This property leads to the trap and store of photons thermalized around the Ly$\\alpha$ frequency for a long time after the cease of the central source emission.  This paper is organized as follows. Section 2 presents the theoretical background of the Ly$\\alpha$ photon transfer in an optically thick medium. Section 3 gives the solution of Ly$\\alpha$ photons emergent from an optically thick spherical halos with a steady source located at the center of the halo. Section 4 is on the solutions when the central source is a flash. A discussion and conclusion are given in Section 5. The details of the numerical implementation is given in the Appendix. ",
        "conclusions": "The resonant scattering made the transfers of resonant photons in physical space and frequency space to be coupled from each other. It leads to the time scale $\\eta$ of the spatial transfer of the resonant photons in the halo with optical depth $\\tau \\gg 1$ being much faster than a purely Brownian diffusion process requiring $\\eta \\propto \\tau^2$. However, essentially it is still a diffusion process, which can be approximately described by $\\eta \\propto \\tau^{1/H}$, with the index $H$ less than but very close to 1. It is possible, if we consider the single longest excursion playing the role of a long-range dependence, that this diffusion process has a positive correlation, or is a fractal Brownian diffusion (e.g. Beran 1994).  The number of photons basically is conserved.  Thus, an optically thick halo is a store of photons with frequency $\\sim \\nu_0$. The time scale of the store is the same as the time scale of the above-mentioned diffusion, i.e. approximately proportional to the optical depth $\\tau$. Moreover, the stored photons are always in the state of local Boltzmann distribution, even when the mean intensity is highly time-dependent. The initial conditions are forgotten in the process of approaching the locally thermal equilibrium. The local Boltzmann distribution is independent of the frequency spectrum and time-dependence of the source.  All these features show that the $\\nu \\simeq \\nu_0$ photons play the central role of radiation transfer with resonant scattering. The major difference between our solutions and some analytical solutions (Harrington 1973; Neufeld 1990; Dijkstra et al. 2006) is also around $x =0$. Since the analytical solutions are based on  the assumption $\\phi(x)=a/\\sqrt{\\pi}x^2$, and the Gaussian core $e^{-x^2}$ is ignored, it generally leads to $J(x=0)=0$. With these approximations the solutions cannot show the effects of restoration and thermalization of photons around $x=0$.  These basic properties found in our numerical solutions yield the following features of Ly$\\alpha$ photons emergent from an optically thick halo.  \\begin{enumerate}  \\item At a given $r$, the profile of the two peaks of the flux is time-independent, and also independent of the initial profile of the photons. Therefore, it is impossible to estimate the line profile of the source from the profile of the flux emergent from an optically thick halo.  \\item The frequencies $|x_{\\pm}|$ of the two peaks of the flux is not less than about 2. This would be useful to estimate the kinetic temperature of the neutral hydrogen atoms.  \\item The resonant scattering makes the flux of the red damping wing very different from that of the DLA model. The flux is non-zero at frequency $\\nu_0$ or $x=0$.  These results would be useful to discriminate the DLA model with models which consider the effect of resonant scattering.  \\item The time scales of light curves of the delayed emission of Ly$\\alpha$ photons from a flash source is mainly determined by the optical depth of the halos. On the other hand, the halo is transparent for high energy photons. A comparison between the light curves of Ly$\\alpha$ photons and high energy photons would be useful to detect the halo.  \\end{enumerate}  For halos with large optical depth, the parameter $\\gamma$ is small even at high redshift. When $\\gamma \\sim 10^{-5}$, the effect of cosmic expansion on photon evolution in the frequency space actually is negligible. All solutions $f(\\eta, r, x)$ given in this paper are almost independent of the Hubble expansion. The effect of $\\gamma$ would be important if the considered time scale $\\eta$ of $f$ is comparable with that of the Hubble expansion."
    },
    "1005/1005.4611_arXiv.txt": {
        "abstract": "The relevant energy ranges for stellar nuclear reactions are introduced. Low-energy compound and direct reactions are discussed. Stellar modifications of the cross sections are presented. Implications for experiments are outlined. ",
        "introduction": "Astrophysical reaction rates describe the number of reactions in a given volume of the stellar plasma and are the essential ingredient for calculations following nucleosynthesis and energy generation in astrophysical environments. They are derived by folding experimental or theoretical reaction cross sections with the energy distribution appropriate for the interacting plasma constituents. For nucleons and nuclei in stars this is the Maxwell-Boltzmann distribution (MBD). A direct determination of the cross sections faces several problems. The majority of nuclei participating in explosive nucleosynthesis is short-lived and thus the reactions cannot (yet) be studied in the laboratory. Thermonuclear reactions on stable and unstable targets proceed at low, often subCoulomb, energies, resulting in tiny reaction cross sections which are further obstacles to experiments. On the theoretical side, several reaction mechanisms may contribute, each with its own peculiar uncertainties. This makes the determination of reaction rates at astrophysically relevant energies a complex puzzle which currently can only be tackled by using several different approaches. In the following I briefly outline a few aspects important for all investigations. ",
        "conclusions": ""
    },
    "1005/1005.3031_arXiv.txt": {
        "abstract": "We have analyzed available full-disc data from the Michelson Doppler Imager (MDI) on board {\\it SoHO} using the ``ring diagram'' technique to determine the behavior of solar meridional flows over solar cycle 23 in the outer 2\\% of the solar radius. We find that the dominant component of meridional flows during solar maximum was much lower than that during the minima at the beginning of cycles 23 and 24. There were differences in the flow velocities even between the two minima. The meridional flows show a migrating pattern with higher-velocity flows migrating towards the equator as activity increases. Additionally, we find that the migrating pattern of the meridional flow matches those of sunspot butterfly diagram and the zonal flows in the shallow layers. A high latitude band in meridional flow appears around 2004, well before the current activity minimum. A Legendre polynomial decomposition of the meridional flows shows that the latitudinal pattern of the flow was also different during the maximum as compared to that during the two minima. The different components of the flow have different time-dependences, and the dependence is different at different depths. ",
        "introduction": "\\label{sec:intro}   Meridional flow plays a key role in the operation of solar dynamo.  Even as early as the 1970s, these flows were  recognized as important in explaining the solar dynamo  and differential rotation (e.g., Durney 1975; Bogart \\& Gierasch 1979). It is believed to be responsible for transporting magnetic elements towards the poles (Wang et al.~1989). It also plays a role in generation of differential rotation and its temporal variations (Glatzmaier \\& Gilman 1982; Spruit 2003). The meridional flows are also invoked as a means to transport flux towards the equator near the base of the convection zone in some solar dynamo models (e.g., Choudhuri et al.~1995; Dikpati \\& Charbonneau 1999). There are no measurements of the meridional flow in the deeper layers of the Sun and hence, near-surface measurements, along with other assumptions,  can be  used to constrain solar dynamo theories.  There is a rich history of the study of the meridional component of flows on the solar surface. (e.g., Schr\\\"oter \\& W\\\"ohl 1975; Howard 1979; Topka et al.~1982; Ribes et al.~1985). However, because the  meridional-flow speed (about 20 m s$^{-1}$) is more than an order of magnitude smaller than other flows at the solar surface, the early observations did not give a clear picture of these flows. Thanks to the improvements in observational techniques, the meridional flow has been reliably measured at the solar surface using both magnetic elements as tracers (Komm et al.~1993; Hathaway \\& Rightmire 2010) and Doppler velocities (e.g., Hathaway 1996). These observations show a general flow towards the poles in both hemispheres that is roughly antisymmetric about the equator.  The development of local helioseismic techniques enabled the study of meridional flow inside the Sun. Using the time-distance technique Giles et al.~(1997) showed that meridional flow persists in the subsurface layers. Schou \\& Bogart (1998), Basu et al.~(1999), Haber et al.~(2000) and others applied the ring diagram technique to study the meridional flow in the subsurface layers up to a depth of about 15 Mm below the surface. The local helioseismic techniques used in these studies do not give any information about the deeper layers. In particular, these could not study the layers close to the base of the solar convection zone which is believed to be the region of the return flow from the poles to the equator. On the other hand, global modes that  can give information about deeper layers are rather insensitive to meridional flow (Roth \\& Stix 2008; Chatterjee \\& Antia 2009) and hence are not useful for this purpose. Nevertheless, even the behavior of meridional flow in the immediate sub-surface layers can  give useful information to constrain  dynamo models. Some models of zonal flows (Spruit 2003) predict a steep variation with depth in the meridional flow near the surface and hence these models too can be tested with the available seismic data.  The availability of seismic data over a full solar cycle has made it also possible to study the time variations of the meridional flows. There have been such studies with data that cover only a part of the solar cycle (Chou \\& Dai 2001; Haber et al.~2002; Basu \\& Antia 2003; Gonz\\'alez Hern\\'andez et al.~2006, 2008, 2010; Zaatri et al.~2006; etc.). All these studies have found significant temporal variations in the meridional flows, though they are not always in complete agreement with each other (see Antia \\& Basu 2007). Among the results that stand out is that the meridional flow speed is slower at maximum activity than at the minimum (Chou \\& Dai~2001; subsequently confirmed by Hathaway \\& Rightmire~2010). Additionally the  pattern of the meridional flow also shows bands of fast or slow flows similar to well known zonal flows (Beck et al.~2002; Gonz\\'alez Hern\\'andez et al.~2010).  In this work we use data  from the Michelson Doppler Imager (MDI) on board {\\it SoHO} to study how the sub-surface meridional flows changed over solar cycle 23. We use the ring diagram technique to study the temporal, latitudinal and depth variations in meridional flow. We chose to use MDI data since it samples the entire solar cycle and hence, may give a clear picture of the temporal variations. It may be noted that although MDI instrument has been functioning through the entire cycle 23, except for a short break around 1998, the ring diagram data are available only during the so-called ``dynamics runs'', which typically cover only 2-3 months every year, earlier during the cycle and for shorter periods since 2003 when {\\it SoHo}'s high-gain antenna started malfunctioning. Thus although the entire solar cycle is sampled by the ring diagram data, the coverage is not continuous and the fill factor is rather low. Data from the Global Oscillation Network Group (GONG) are available for the declining phase of cycle 23 only. These have been analyzed by Gonz\\'alez Hern\\'andez et al.~(2010). GONG data has the advantage of continuous coverage in time from July 2001 onwards. MDI magnetic field data have been used by Hathaway \\& Rightmire (2010) to study surface meridional flows. ",
        "conclusions": "\\label{sec:disc}   We have studied solar meridional flows and their variations from June 1996 to April 2009 using the ring diagram technique applied to full-disc Dopplergrams from MDI. This period covers the  solar cycle 23 in its entirety. The striking feature that we find is that the flows not only vary with depth, and latitude, but that the time-dependence of the variation  is a function of both depth and latitude.  The meridional flows show a migrating pattern, very similar to that of the zonal flows. This result was first seen in time-distance measurements of Beck et al.~(2002) and they had concluded that there is a strong correlation between the time-dependent part of the meridional flows and the torsional oscillations (zonal flows). This was later confirmed using GONG data by Gonz\\'alez~Hern\\'andez et al.~(2010). However, Beck et al.~(2002) only had data for the rising part of the cycle, while Gonz\\'alez~Hern\\'andez et al.~(2010), using GONG data, could only study the declining phase of cycle 23. As a result, it was not completely clear whether the meridional flow pattern matches the zonal flow bands. Our results confirm the migrating pattern, and also that close to the surface the migration mimics the migration of the zonal flows. We find that a band of faster flow at high latitudes has appeared around 2004, long before the current minimum, this is also seen with GONG data (Gonz\\'alez~Hern\\'andez et al.~2010). However, we find that although the pattern of the zonal flow is similar from the surface down to a radius of $0.90R_\\odot$ (e.g. Antia et al.~2008), the meridional flow pattern changes rapidly with depth. It is possible that flow pattern at the deeper levels reflects the magnetic field pattern at those levels, however, the question then would be the lack of the change in the zonal flows. It is, however, not completely clear if the differences in the depth-dependence of the zonal and meridional flow patterns is of a fundamental nature, or  simply a results of the fact that ring diagram results have a higher radial resolution near the solar surface compared to global mode analyses used to study zonal flows.   Svanda et al.~(2008) have studied the effects of solar active regions on meridional flows and claim that a part of the temporal variation at low latitudes could be explained by flows around active regions. This could explain the shifting pattern of the meridional flows, but not all of the observed  temporal variation. However, by  masking out active regions  Gonz\\'alez~Hern\\'andez et al.~(2008) showed that while the resultant flow velocity that they obtain changes in the process, there is a second component that still shows a decreasing amplitude as magnetic activity increases. Thus what we are seeing is not merely a local effect of sunspot magnetic fields. In addition to sunspots, there is some evidence to suggest that the variation in the B-angle, $B_0$. with time also may cause apparent temporal variation in meridional flows (Zaatri et al.~2006). This has been argued to be the reason why some results showed presence of counter cells in layers below the surface (Haber et al.~2002; Gonz\\'alez Hern\\'andez et al.~2006). However, in this work we have used only the antisymmetric component of the meridional flow, which minimizes the effect of $B_0$ angle variation. Thus we do not expect our results to be significantly affected by variation in $B_0$. The fact that we see a consistent variation in the meridional-flow velocity supports this conclusion.    In order to get a better understanding of the flow pattern, we decompose the latitudinal dependence in terms of associated Legendre polynomials. This decomposition of meridional flows shows that higher degree components also make a significant contribution. These components do not have the same sign at all latitudes in a given hemisphere, and hence, they give rise to a complex pattern of variation in the resulting flows. These components, in particular, the $P^1_6$ and $P^1_8$ components, change sign with time and are reversed during the solar maximum with respect to the flow during the solar minimum. While the higher order components show a time-variation, the changes are not clearly correlated to solar activity. This is probably the result of the fact that global activity indices are obtained by averaging over the disc while the higher components of flows have a complex latitudinal dependence.  The dominant component of meridional flow ($P_2^1(\\cos\\theta)$) shows a clear variation with time that is anti-correlated with global solar activity. Chou \\& Dai (2001) also found similar results. Basu \\& Antia (2003) had also found similar variations but the result was not very clear with the limited data available at that time. The extended minimum before cycle 24 clearly shows up in the variation of coefficient $C_2$ in the shallower layers. There is no sign until the last data set (June 2009) of $C_2$ starting to decrease. Our results indicate that the magnitude of this component is similar between the two minima of cycles 23 and 24. Our results are qualitatively similar to those found by Hathaway \\& Rightmire (2010) from magnetic-field observations made by MDI, though the amplitude they find is consistently lower than what we find after correcting for the difference in normalization. We get $C_2$ values of $12.8\\pm 0.4$ m s$^{-1}$, $6.9 \\pm 0.4$  m s$^{-1}$ and $12.5 \\pm 0.5$ m s$^{-1}$ respectively, for mid-1996, early 2002 and mid 2009 at a depth of 1.4 Mm ($r=0.998R_\\odot$). It should be noted that  Hathaway \\& Rightmire (2010) report the amplitude of the $2\\sin\\theta\\cos\\theta$ term while we use $P_2^1(\\cos\\theta)=-3\\sin\\theta\\cos\\theta$. Thus our $C_2$ should be multiplied by 1.5 before comparing with their values. Hathaway \\& Rightmire (2010) also find the coefficient $C_2$ to be larger during the current minimum as compared to that during the previous minimum. We however, do not find a significant difference  at the lowest activity levels that we have access to, however $C_2$ during slightly higher activity levels (albeit still low activity) are somewhat different during the minimum of cycle 24 compared with those of cycle 23. The differences are clearly seen in the coefficients of the higher order components in Legendre decomposition. As can be seen from Table~1, the activity level during the first set in our list is the same as that during the last set. Unfortunately, we do not have data for December 2008, when solar activity was lowest. It is quite possible that the difference between our results and those of Hathaway \\& Rightmire (2010) is because we are looking at different depths, or because of the different observables and techniques used in the two studies. Our results show a rapid variation of flow amplitudes with depth and it is  not clear what depth is probed by the magnetic features used by Hathaway \\& Rightmire (2010).   The strength of the solar meridional flow is believed to play an important role in transporting magnetic flux. Devore et al.~(1984) show that the meridional flows play a significant role in magnetic field distribution at low activity. A lower meridional flow speed tends to give stronger polar fields (Sheeley et al. 1989). Our results (see Figure~\\ref{fig:merid4}) would imply a weaker polar field during the minimum of cycle 24. That is indeed the case. Wang et al.~(2009) find that the  polar fields during the minimum of cycle 24 were about 40\\% weeker than that during the earlier minimum.  As mentioned earlier, we find that the meridional flow speed increases with depth. The increase in meridional flow amplitudes with depth in the near-surface layers is seen in {\\it ad hoc} models of the solar  meridional flow (e.g., Roth et al.~2002; Chatterjee \\& Antia 2009). In these models the increase can be attributed to the boundary condition of radial velocity vanishing at the top boundary and the small density scale heights in the subsurface layers. In these models the meridional velocity has a maximum close to the surface, around a depth of 3.5 Mm ($r=0.995R_\\odot$), and then decreases becoming significantly less below a depth of 14 Mm ($r=0.98R_\\odot$). Since these features are essentially a result of the continuity equation, similar variations may be seen in more realistic profiles or in the Sun. On the other hand, Spruit (2003) has proposed a model for torsional oscillations as a geostrophic flow due to lower subsurface temperatures in active regions. This model also gives rise to a meridional flow. However, it predicts that the meridional flow amplitude would decrease with depth, which is opposite to what we find. If this trend is confirmed, such models will be ruled out.  Our results are for very shallow layers of the Sun; we can only probe the outer 2\\% of solar radius with any reliability. There would be considerable scientific return if the analysis could be extended to deeper layers not covered in this work. Gonz\\'alez Hern\\'andez et al~(2006) have tried getting results for deeper layers of the Sun using ring diagrams of regions of size 32\\dg\\ (in contrast, this and other ring diagram work use 16\\dg\\ regions) and find they can cover the outer 5\\% of the Sun. However, this technique has the drawback that it necessarily degrades the latitudinal resolution. Perhaps a combination of rings-diagrams of different sizes is required to get a good coverage. Time-distance helioseismology could, in principle, give us information of flows in the deeper layer, but again there would be a tradeoff in terms of latitudinal resolution. Data from the Helioseismic and Magnetic Imager (HMI) on board the recently launched Solar Dynamics Observatory will undoubtedly provide us with a better picture of solar meridional flows, though analyzing meridional flows in the deeper layers will remain a problem. The data, however, will allow us to study cycle-to-cycle variations of the flows.  The higher resolution of HMI will also allow us to extend helioseismic meridional flow studies to higher latitudes."
    },
    "1005/1005.1943_arXiv.txt": {
        "abstract": "{We present \\herschel\\ observations of the isolated, low--mass star--forming Bok globule CB244.  It contains two cold sources, a low--mass Class 0 protostar and a starless core, which is likely to be prestellar in nature, separated by 90$\\arcsec$ ($\\sim\\!18000$~AU).  The \\herschel\\ data sample the peak of the Planck spectrum for these sources, and are therefore ideal for dust--temperature and column density modeling.  With these data and a near-IR extinction map, the MIPS 70~\\micron\\ mosaic, the SCUBA 850~\\micron\\ map, and the IRAM 1.3~mm map, we model the dust--temperature and column density of \\cb\\ and present the first measured dust--temperature map of an entire star--forming molecular cloud.  We find that the column--averaged dust--temperature near the protostar is $\\sim\\!17.7$~K, while for the starless core it is $\\sim\\!10.6$~K, and that the effect of external heating causes the cloud dust--temperature to rise to $\\sim\\!17$~K where the hydrogen column density drops below $10^{21}$~cm$^{-2}$.  The total hydrogen mass of CB\\,244 (assuming a distance of 200~pc) is $15\\pm5$~\\msun. The mass of the protostellar core is $1.6\\pm0.1$~\\msun\\ and the mass of the starless core is $5\\pm2$~\\msun, indicating that $\\sim$45\\% of the mass in the globule is participating in the star--formation process.} ",
        "introduction": "Bok globules are nearby, small, relatively isolated molecular clouds undergoing low--mass star--formation \\citep[e.g.,][]{clemens88,laun97}.  They tend to have only one or two star--forming cores that are embedded in a larger common cloud.  These relatively simple characteristics make these objects ideal for studying the detailed processes taking place in low-mass star formation.  Specifically, the temperature and density structure are fundamental physical parameters necessary to understand core fragmentation, collapse, and chemical evolution \\citep[e.g.,][]{wt07,stutz08,stutz09,laun10}.  The {\\it Herschel} \\citep{pil10} data cover the wavelength range which samples the peak of the Planck spectrum for cold sources (6 - 20~K).  This wavelength regime is critical for accurate modeling of the temperature and density structure in the cold environments where stars are born \\citep[e.g.,][]{shetty09b}.  The {\\it Herschel} Guaranteed Time Key Program ``Earliest Phases of Star--formation'' (EPoS; P.I.\\ O. Krause) sample consists of the Photodetector Array Camera and Spectrometer \\citep[PACS;][]{pog10} and the Spectral and Photometric Imaging Receiver \\citep[SPIRE;][]{griff10} imaging--mode observations of sites of both high-- and low--mass star formation.  The low--mass Science Demonstration Phase portion of the sample is the focus of this contribution.  The source \\cb\\ \\citep[L1262;][]{lynds62} is a Bok globule at a distance of $\\sim\\!200$~pc \\citep{hilton95}, with an approximate extent of $\\sim\\!6\\arcmin$, or about 0.5~pc.  The \\cb\\ globule contains two submm {peaks, one associated with a Class 0 protostar located at R.A.\\ = $23^h 25^m 46.3^s$, Decl.\\ = $+74\\degr 17\\arcmin 39.1\\arcsec$, and one associated with a starless core located at R.A.\\ = $23^h 25^m 27.1^s$, Decl.\\ = $+74\\degr 18\\arcmin 25.3\\arcsec$.  The protostar drives a molecular outflow \\citep[e.g.,][]{clemens91}, while the detection of the \\cb\\ starless core was first published as an additional source inside the globule by \\citet{laun96} and \\citet{shirley00} and produces both an 8~\\micron\\ \\citep{tobin10} and a 24~\\micron\\ shadow.  The YSO and starless core are separated by $\\sim\\!90\\arcsec$, and are therefore well resolved throughout the PACS and SPIRE bands.  We use \\herschel\\ imaging to construct spatially resolved spectral energy distributions (SEDs) of the entire cloud that cover both sides of the peak of the SED. These data allow us to reconstruct the dust--temperature map and column density distribution of the Bok globule and hence the density profiles and mass distribution with unprecedented accuracy. They also reveal the role of external heating and shielding by the envelopes in the energy balance of such cores in isolated globules.  \\begin{figure*} \\begin{center} \\scalebox{0.7}{{\\includegraphics{f1.eps}}} \\caption{$8\\arcmin \\times 8\\arcmin$ images of \\cb, shown at the wavelengths indicated in the bottom--right of each panel.  The images are shown on a log scale at stretches intended to highlight the structure of absorption or emission at each wavelength.  The locations of the protostar and starless core, derived from the 8~\\micron\\ image, are indicated with $\\times$--symbols (note the offset between the $\\times$'s and the starless core emission at 100 -- 250~\\micron\\ wavelengths).  The panel contours `are labeled with their corresponding wavelengths.  The contour levels are: $A_V$ extinction map levels = \\{5,10,15,20,25\\}~mags (inner dotted contours indicates the region inside which we place lower--limits on the extinction); SPIRE 350~\\micron\\ = \\{0.5,1.0,1.5,3.5\\}~mJy/$\\sq\\arcsec$; SCUBA 850~\\micron\\ = \\{0.3,0.75\\}~mJy/$\\sq\\arcsec$; IRAM 1.3~mm = \\{0.1,0.3\\}~mJy/$\\sq\\arcsec$.  Approximate beam sizes are indicated as yellow circles. \\label{fig:img}} \\end{center} \\end{figure*} ",
        "conclusions": "As shown in Fig.~\\ref{fig:temp}, the line--of--sight--averaged dust--temperature decreases constantly without any significant jump from $\\sim\\!17\\pm1$~K at the outer boundary of the cloud, where $N_{\\rm H}$ drops below $10^{21}$\\,cm$^{-2}$\\ and $A_V \\approx 0.3$\\,mag, to $\\sim\\!13$\\,K in the inner cloud at $N_{\\rm H}\\approx 1$ --- $2 \\times 10^{22}$\\,cm$^{-2}$\\ and $A_V \\approx 3$ -- $5$\\,mag. Note that due to the two major drawbacks of our simplified approach, beam--smoothing and line--of--sight--averaging, the local dust--temperatures at the source centers will be higher in the protostar and lower in the prestellar core.  Nevertheless, the boundaries of the two sources can be derived from the dust--temperature map rather than from the smooth column density map. We find that $T_{\\rm d } = 13$\\,K marks the mean inner dust--temperature of the globule, from where the temperature towards the protostar rises steeply and that of the prestellar core decreases. We therefore adopt ellipses that follow the 13\\,K contour in the dust--temperature map as boundaries for the two sources (see Fig.~2). The mean radii of these ellipses are 50\\arcsec\\ (10,000~AU) for the protostellar core and 76\\arcsec\\ (15,000~AU) for the prestellar core. Integration of the column density map within these boundaries yields $M_{\\rm H} = 1.6\\pm0.3$\\,M$_{\\odot}$\\ for the protostellar core and $M_{\\rm H} = 5\\pm2$\\,M$_{\\odot}$\\ for the prestellar core, where the uncertainties are derived from an assumed $\\pm 1$\\,K uncertainty in the dust--temperature.  The total mass of the globule (within the $N_{\\rm H} = 1\\times 10^{21}$\\,cm$^{-2}$\\ contour) is $M_{\\rm H}=15\\pm5$\\,M$_{\\odot}$.  We note that $M_{\\rm gas} = 1.36 M_{\\rm H}$.  These masses imply that $\\sim$45\\% of the mass of the globule is participating in the star--formation process.  Including the new \\herschel\\ data, we derive the following parameters for the protostar: L$_{\\rm bol}\\!\\sim\\!1.5$~\\lsun, L$_{\\rm submm}$/L$_{\\rm bol}\\!\\sim 4$\\%, T$_{\\rm bol}\\!\\sim\\!62$~K, and M$_{\\rm env}$/L$_{\\rm bol}\\!\\sim\\!1.45$~\\msun/\\lsun, \\citep[cf.][]{laun10}, confirming the Class\\,0 classification according to \\citet{chen95} and \\citet{andre00}.  Furthermore, the C$^{18}$O (2--1) FWHM line width for the prestellar core is $\\sim\\!0.9$~km~s$^{-1}$ (Stutz et al., in prep.); using our fitted prestellar core mass and temperature, we find that the ratio of gravitational energy to thermal and turbulent energy is E$_{\\rm grav}$ / (E$_{\\rm therm}$ $+$ E$_{\\rm turb}) \\simeq 0.9$, a marginally sub--critical value which is expected for prestellar cores.  For comparison, \\citet{tobin10} derive a mass of $\\sim$3 -- 4~\\msun\\ for the prestellar core (integrated over a similar area and scaled to a distance of 200~pc) using the 8~\\micron\\ shadow (see Fig.~1), a good agreement given the uncertainties in both mass derivation methods.  The prominent and extended 3.6~\\micron\\ coreshine (see Fig.~1), originating from dust--grain scattering of the background radiation field, is an indication of grain growth \\citep{stein10} in the \\cb\\ cloud.  These pieces of evidence together indicate that the \\cb\\ globule, and other globules like it, are excellent sources in which to study the earliest phases of low--mass star--formation.  As a next step and in a follow--up paper, we will employ 3D--modeling to overcome the effects of line--of--sight averaging and beam-smoothing, and to reconstruct the full dust--temperature and density structure of \\cb\\ and the other sources in our sample.  Measuring reliable temperatures and column densities with the \\herschel\\ data in a sample of prestellar and protostellar cores is a fundamental step towards revealing the initial conditions of low--mass star--formation."
    },
    "1005/1005.0174_arXiv.txt": {
        "abstract": "We present an analysis of the science image quality obtained on the twin 6.5~m Magellan telescopes over a 1.5 year period, using images of $\\sim10^5$ stars. We find that the telescopes generally obtain significantly better image quality than the DIMM-measured seeing. This is qualitatively consistent with expectations for large telescopes, where the wavefront outer scale of the turbulence spectrum plays a significant role. However, the dominant effect is found to be wind speed with Magellan outperforming the DIMMs most markedly when the wind is strongest. Excluding data taken during strong wind conditions ($>10$~m~s$^{-1}$), we find that the Magellan telescopes still significantly outperform the DIMM seeing, and we estimate the site to have $\\mathcal{L}_0\\sim25$~m on average. We also report on the first detection of a negative bias in DIMM data. This is found to occur, as predicted, when the DIMM is affected by certain optical aberrations and the turbulence profile is dominated by the upper layers of the atmosphere. ",
        "introduction": "\\label{sec-intro} Much time and money is invested in finding optimal sites for the next generation of big optical-IR telescopes and it is of interest to learn the detailed atmospheric properties of major observatory sites since the seeing is a major limitation to the resolution and signal to noise ratio that can be achieved. It is well known that large aperture telescopes frequently perform better than we would naively expect from the simple seeing measurements at major observatories (e.g.~\\citealt{martin+98, wilson+99, tok+07,salmon+09,els+09}) in terms of image full width at half maximum (FWHM). Most astronomers consider the ``seeing'' to be synonymous with the FWHM of a stellar image on their detector, but this assessment is formally incorrect. Seeing is a simple measurement of Fried's parameter, $r_0$~\\citep{fried65}, determined using differential image motion monitors (DIMMs -- see~\\citealt{DIMM} and references therein) but there are other important parameters governing imaging quality, such as the characteristic time constant, the isoplanatic angle and the wavefront outer scale, \\Lo. These can play important roles in different regimes.   \\subsection{Background: Turbulence, seeing and outer scale} Using a single parameter to determine image FWHM implicitly assumes that energy is injected into the atmosphere at infinite scale following the Kolmogorov turbulence model~\\citep{tatarskii69,kolm41}, an approximation that is valid only for small telescopes. In the Kolmogorov model, energy is gradually transferred to smaller and smaller scales until it is dissipated through the viscosity of the air on scales of $\\sim$ mm -- cm (the inner scale, $l_0$). The result is local variations in atmospheric temperature and density, and thus in refractive index, $n$. In reality the atmosphere continually has energy injected into it through heating, convection currents and wind shear on scales $L_0\\sim10$~m, which we can introduce as a cutoff in the turbulence spectrum or ``outer scale of turbulence'', \\Lo. This is known as the \\vk turbulence model (see~\\citealt{borgnino90,ziad+00}).  In practice atmospheric turbulence results in small amounts of image motion. The standard device for measuring this motion is the DIMM. A DIMM measures the variance of the differential image motion, $\\sigma_d^2$, of two images of the same star, observed through separate apertures in the entrance pupil of a small telescope. This gives a measurement of \\ro\\ for a sub-aperture of diameter $D$ and wavelength $\\lambda$: $\\sigma_d^2 \\propto \\lambda r_0^{-5/3} D^{-1/3}$. Thus the FWHM $\\epsilon_0$ of the long-exposure seeing-limited point spread function in large telescopes is computed within the Kolmogorov turbulence model as $\\epsilon_0 = 0.98 \\lambda/r_0 \\propto ( D/\\lambda )^{1/5} (\\sigma_d^2)^{3/5}$.   From a turbulence spectrum it is possible to derive a statistical structure function $F_\\phi(\\mathbi{r})$ which describes the phase decorrelation between any pair of points in the atmosphere, separated by $\\mathbi{r}$. % In the Kolmogorov model, the phase structure function (the mean square differential phase error -- see~\\citealt{roddier81} for a useful derivation) grows larger with distance of separation, $D_\\phi(\\mathbi{r}) \\propto (|\\mathbi{r}|/r_0)^{5/3}$  -- see dotted line in Fig.~\\ref{fig-strucfun}. But in the \\vk model the structure function becomes more complicated~\\citep{ConRonch72} and grows rapidly with $r$ up to $r = \\mathcal{L}_0$ where it saturates -- see solid line in Fig.~\\ref{fig-strucfun}: Beyond \\Lo, $D_\\phi(\\mathbi{r})$ no longer increases. The outer scale effect is detectable in telescopes that are an appreciable fraction of the size of \\Lo, where it begins to deviate from the Kolmogorov case.  \\begin{figure}[htbp] \\begin{center} \\plotone{f1.eps}% \\caption{\\label{fig-strucfun} The \\vk phase structure function, $D_\\phi(r)$ (solid line $\\mathcal{L}_0=25$~m, dashed line $\\mathcal{L}_0=1000$~m) compared to the Kolmogorov phase structure function (dotted line), all for $r_0=0.2$~m. Unlike the Kolmogorov structure function which continues to grow with increasing $r$, the \\vk structure function saturates at $r=$\\Lo. The outer scale effect becomes noticeable at scales significantly smaller than \\Lo.} \\end{center} \\end{figure}  From the structure function we can derive the optical transfer function (OTF) of a large telescope, $T_0(\\mathbi{f}) = \\exp[-0.5 D_\\phi (\\lambda\\mathbi{f})]$, where $\\mathbi{f}$ is the spatial frequency in inverse radians. % The corresponding Point Spread Function (PSF) is found by taking the Fourier Transform of the OTF. % \\citet{tok02} obtained the PSF for the \\vk case by renormalizing the OTF. % The result is an approximation of the seeing-limited PSF of a large telescope, with FWHM $\\epsilon_{\\mathrm vK}$ approximated by: \\begin{equation} \\label{eq:vk} \\left( \\frac{\\epsilon_{\\mathrm vK}}{\\epsilon_0} \\right)^{2} \\approx 1 - 2.183 \\left( \\frac{r_0}{\\mathcal{L}_0} \\right)^{0.356} \\end{equation} which we use later in this paper to estimate an average value of \\Lo\\ at the Las Campanas site. The effect of introducing an outer scale to the phase structure function is to decrease the FWHM of the point spread function (PSF) of a ``large'' telescope. ",
        "conclusions": "We have presented a simple method for measuring the outer scale of turbulence that can be used at any observatory site that has a large aperture telescope used regularly for imaging, and a DIMM device for measuring the seeing. The Las Campanas Observatory site offers truly excellent seeing conditions. The effective ``seeing'' for the Magellan telescopes is considerably underestimated by the DIMMs. While the DIMMs give a median seeing of 0\\farcs625, the median Magellan science image FWHM is 0\\farcs575. We find the dominant cause of this difference is local to the DIMM: high wind velocities introduce local turbulence, perhaps through interaction with the clamshell dome. There is weak evidence for a directional influence on the DIMM image degradation at wind directions perpendicular to the orientation of the opening. However, removing points with wind speeds $>10$~ms$^{-1}$ we find that the Magellan image quality is still significantly better than the DIMM seeing. From the cumulative distributions of the DIPSF and DIMM FWHMs we estimate that the site has an average turbulence outer scale of $\\mathcal{L}_0\\sim25$~m, similar to that measured at other sites.  We were unable to include our 2009 data due to a mis-collimation of the Cerro Manqui DIMM. This leads to both positive and negative biases in the DIMM data for the year, and is the first time that such negative biases have been recorded, although they are predicted. These negative biases occur when the upper layers of the atmosphere dominate the atmospheric turbulence profile."
    },
    "1005/1005.2989_arXiv.txt": {
        "abstract": "The Szekeres inhomogeneous models can be used to model the true lumpy universe that we observe. This family of exact solutions to Einstein's equations was originally derived with a general metric that has no symmetries. In this work, we develop and use a framework to integrate the angular diameter and luminosity distances in the general Szekeres models. We use the affine null geodesic equations in order to derive a set of first-order ordinary differential equations that can be integrated numerically to calculate the partial derivatives of the null vector components. These equations allow the integration in all generality of the distances in the Szekeres models and some examples are given. The redshift is determined from simultaneous integration of the null geodesic equations. This work does not assume spherical or axial symmetry, and the results will be useful for comparisons of the general Szekeres inhomogeneous models to current and future cosmological data. ",
        "introduction": "Recently, there has been a renewed interest in studying some current questions in cosmology using inhomogeneous models, see for example the reviews \\cite{Celerier2007,BolejkoEtAlBook} and references therein. Studies have included for example the formation of structures in the universe \\cite{KandH,HandK,Bolejko2006a,Bolejko2007} and the question of cosmic acceleration, e.g. see \\cite{ Iguchietal2002,Alnesetal2006,Enqvist1,Garfinkle2006,Kaietal2006,Biswasetal2006,TanimotoNambu2007,Hunt,Enqvist,IshakEtAl2008,Chung,Garcia} and references there. It is also of interest that such inhomogeneous models can offer a wider range of interpretation of cosmological observations.  Whereas one can find a large body of literature studying exact solutions to Einstein's field equations that are inhomogeneous cosmological models \\cite{krasinski1997,KSHM}, relatively little work has been done comparing inhomogeneous models to various observations.   In this paper, we consider the calculation of the area and luminosity distances in the general Szekeres models. The affinely parameterized null geodesic equations are used to derive a set of first-order ordinary differential equations ready to be integrated numerically to calculate the partial derivatives of the null vector components. The results allow the integration of the distances in the general Szekeres models. The redshift is   calculated from numerical integration of the set of null geodesics. The general results should be of immediate application to comparisons between the general Szekeres models and cosmological data.  The models were originally derived by Szekeres \\cite{Szekeres1,Szekeres2} as an exact solution to Einstein's equations with a general metric that has no symmetries (i.e. no Killing vector fields) and a dust source. The models have been investigated analytically by several authors \\cite{Szafron1977,SussmanAndTriginer1999,BonnorAndTomimura1976,Bonnoretal1977,GoodeAndWainwright1982,Bonnor1985,Barrow,SussmanAndTriginer1999,HellabyAndKrasinski2002,Bolejko2006a,Bolejko2007,HellabyAndKrasinski2008,Krasinski2008,Nolan2007}. Reference \\cite{EllisVanEllst1998} classifies them in the same category as the observed lumpy universe (see the table on page 37 there) and the models are considered as one of the best exact solution candidates to represent the true lumpy universe we live in. The models have been discussed in regards to structure formations in \\cite{GoodeAndWainwright1982,Bolejko2006a,Bolejko2007} and also most recently in the context of cosmic acceleration and comparisons to supernova data in \\cite{IshakEtAl2008,BolejkoCelerier} and the Rees-Sciama effect in the cosmic microwave background radiation in \\cite{Bolejko2009}.  It is perhaps worth clarifying that we are not proposing the Szekeres model as the true and ultimate model of the universe but rather developing a framework based on exact inhomogeneous cosmological models where cosmological observations can be investigated with a wider range of possible interpretations. ",
        "conclusions": "We derived and used a framework to integrate the area and luminosity distances in the general Szekeres models. We used the general affinely parameterized null geodesic equations in order to derive a set of first-order ordinary differential equations that can be integrated numerically to calculate the partial derivatives of the null vector components. These equations allow the numerical integration of the area and luminosity distances in the general Szekeres models. We determined the redshift from simultaneous integration of the null geodesic equations. This work does not assume spherical or axial symmetry and will be useful for comparisons of the general Szekeres inhomogeneous models to current and future cosmological data."
    },
    "1005/1005.0032_arXiv.txt": {
        "abstract": "We report the discovery of very high energy gamma-ray emission from the direction of the SNR G54.1+0.3 using the VERITAS ground-based gamma-ray observatory.~The TeV signal has an overall significance of 6.8$\\sigma$ and appears point-like given the % resolution of the instrument.~The integral flux above 1 TeV is 2.5\\% of the Crab Nebula flux and significant emission is measured between 250 GeV and 4 TeV, well described by a power-law energy spectrum dN/dE $\\sim$ E$^{-\\Gamma}$ with a photon index $\\Gamma= 2.39\\pm0.23_{stat}\\pm0.30_{sys}$.~We find no evidence of time variability among observations spanning almost two years.~Based on the location, the morphology, the measured spectrum, the lack of variability and a comparison with similar systems previously detected in the TeV band, the most likely counterpart of this new VHE gamma-ray source is the PWN in the SNR G54.1+0.3.~The measured X-ray to VHE gamma-ray luminosity ratio is the lowest among all the nebulae supposedly driven by young rotation-powered pulsars, which could indicate a particle-dominated PWN. ",
        "introduction": "In recent years, new insights into the most extreme environments in the Galaxy have been obtained through the very high energy~(VHE, $>$100 GeV)~gamma-ray band~\\citep{aharonian05a}.~The majority of the $>$50 Galactic TeV sources found thus far are associated with supernova remnants (SNRs) or pulsar wind nebulae (PWNe)\\footnote{TeVCat, an online catalog for TeV Astronomy, http://tevcat.uchicago.edu/}.~The detection of these stellar remnants in VHE gamma rays constitutes a probe of non-thermal astrophysical processes.~The acceleration of  electrons or nuclei, above 1 TeV,  in these sources leads to the generation of gamma rays.~Relativistic electrons can produce gamma rays by non-thermal bremsstrahlung, and by inverse-Compton (IC) scattering on ambient microwave, IR, or optical photons, whereas protons and atomic nuclei can generate gamma rays via the decay of neutral pions produced in hadronic interactions with ambient material.~A confirmation of the latter process associated with the gamma rays from SNRs is seen as a probe of nucleonic cosmic ray production sites in the Galaxy. \\par The SNR G54.1+0.3 has been predicted to emit gamma rays detectable with the current generation of TeV instruments after the discovery of a rotation-powered pulsar in its centre, a source of high energy electrons~(e.g.~\\cite{bednarek05}).~This pulsar, PSR J1930+1852~(P=136 ms), with a derived spindown power of $\\dot{E}$ = 1.2$\\times10^{37}$ erg/s and a spindown age of $\\tau_c$$\\sim$2900 years~\\citep{camilo02}, is among the youngest and most energetic pulsars known.~The PWN powered by this pulsar was first recognized as such from high resolution radio observations~\\citep{reich85}.~The PWN exhibits a filled-center morphology~($2'\\times1.5'$), a flat radio spectrum~($\\alpha\\approx0.13$) which is fairly constant over the entire source extent, and a high level of polarization~\\citep{velusamy88}.~The discovery of an X-ray jet,~consistent with the radio extension,~and the measurement of its non-thermal spectrum confirmed the PWN nature of G54.1+0.3~\\citep{lu01}.~More detailed X-ray observations by~\\cite{lu02} resolved the inner structure of the nebula as a central point-like source with a surrounding ring and bipolar elongations, comparable to the Crab Nebula.~The same authors suggested a distance of $\\sim$5 kpc, based on the X-ray absorption.~The most recent estimate is 6.2$^{+1.0}_{-0.6}$ kpc, based on an apparent morphological association between the PWN and a nearby molecular CO cloud~\\citep{leahy08}. \\par Similar to the Crab Nebula, G54.1+0.3 lacks a strong surrounding circumstellar interaction between the SNR forward shock and the progenitor star wind~\\citep{chevalier05}, in which particles could be accelerated and produce an additional non-thermal emission contribution.~Recently, \\cite{leahy08} found a shell of 30' diameter surrounding G54.1+0.3, but the authors concluded that these are unlikely to be associated based on the different measured distances.~After that,~\\cite{koo08} identified a partial shell-like infrared structure surrounding much of the X-ray/radio core in G54.1+0.3, which they proposed to be a star-forming region triggered by late-phase winds from the SNR progenitor star.~Based on high-resolution IR spectra, however,~\\cite{temim09} suggested that the shell is composed of freshly-formed ejecta dust that appears to be heated by the expanding PWN. \\par G54.1+0.3 has not been previously detected at $\\gamma$-ray energies.~The source lies 1.4$^{\\circ}$ from the EGRET source  3EG J1928+1733~\\citep{hartman99}.~This is slightly outside the 99\\% confidence-level error box.~The older~($\\tau_{c}$$\\sim$8.2$\\times10^{4}$ years)~and less-energetic ($\\dot{E}$ = 1.6$\\times10^{36}$ ergs/s) pulsar PSR J1928+1746 is proposed as a more likely counterpart to the unidentified 3EG source~\\citep{cordes06}.~This is strengthened with the much better localization of the gamma-ray source 1FGL J1929.0+1741c,~associated with 3EG J1928+1733 in the first Fermi catalogue~\\citep{FermiCatalogue}.~This other pulsar is 1.2$^{\\circ}$ away from PSR J1930+1852.~In the VHE $\\gamma$-ray domain, the only published observation is an upper limit from the HEGRA collaboration at $20\\%$ of the Crab Nebula flux above 600 GeV~\\citep{aharonian02}.~In this Letter we report the discovery by VERITAS of VHE gamma-ray emission from the direction of the very young system G54.1+0.3/PSR J1930+1852.~The instrument and observations are described in \\S~2,~while the analysis performed on these data is detailed in \\S~3.~Results are presented in \\S~4 and interpretation and conclusion in \\S~5. ",
        "conclusions": "VERITAS has detected VHE gamma-ray emission from  the direction of the SNR G54.1+0.3, providing a new input to study the non-thermal particle population in this object.~This new TeV source, VER J1930+188,~appears point-like within the instrument resolution and does not show a notable offset from the pulsar PSR J1930+1852, placing it among the more compact TeV sources found to be associated with an SNR.~The emission is co-located with the X-ray and radio PWN in the SNR G54.1+0.3,~and also with a molecular CO cloud in apparent morphological association with the PWN and a partial shell of IR ejecta around the PWN.  \\par Assuming the gamma-ray production in VER J1930+188 is related to the pulsar, the compactness of the TeV source and lack of significant offset between the pulsar and its X-ray PWN and the TeV source can be attributed to its relative youth.~The spin-down age of PSR J1930+1852 is 2900 years, which places it as the latest addition to an emergent  class of TeV sources associated with very young pulsars:  PSR B0531+21/Crab Nebula~\\citep{weekes89}, PSR J1747-2809/HESS J1747-281 in G0.9+0.1~\\citep{camilo09,aharonian05b}, PSR J1845-0258/HESS J1845-029 in SNR Kes 75, PSR J1833-1036/HESS J1833-105 in G21.5-0.9~\\citep{djannati07} and PSR J1813-1749/HESS J1813-178 in G12.82-0.02~\\citep{gotthelf09}.~Those PWNe are good candidates for being in the early stage of evolution, where the wind nebula is interacting with the freely-expanding supernova ejecta.~This population stands in contrast to an apparently more common variety of energetic-pulsar-associated TeV sources, a group which contains, as notable examples, PSR B1823-1313/HESS J1825-137~\\citep{aharonian06a},~PSR B0833-45/Vela X~\\citep{aharonian06b} and PSR J2229+6114/VER J2227+608~\\citep{acciari09}.~These objects show extended TeV emission which is both larger than the X-ray PWN and offset from the pulsar location~(see, e.g, \\cite{deJager08}).~A common element in this class is older, middle-aged pulsars, with ages $>$10$^4$ yr.~Here the PWN may be displaced in some direction, possibly by an asymmetric SNR reverse shock generated from an explosion into an inhomogeneous medium~\\citep{blondin01}. \\par Table 2 lists the pulsar properties (spin-down power, age and distance) and the nebular energetics in several compact PWNe detected in TeV.~The total power, L$_g$, radiated by G54.1+0.3 in the energy band 0.2 - 10 TeV is $2.1\\times10^{34}$ erg s$^{-1}$ compared to $8.6\\times10^{34}$ erg s$^{-1}$ for the Crab Nebula. This value of the gamma-ray luminosity of G54.1+0.3 is  the same order of magnitude what has been found for the other PWNe in the table.~A more relevant quantity to compare is the efficiency of TeV luminosity production compared to the pulsar power which is $\\eta_{\\gamma}\\sim$0.17\\% for G54.1+0.3.~This is similar to that of G0.9+0.1~($\\approx$~0.1\\%) and 10 times larger than other young systems like the Crab Nebula~($\\approx$~0.02\\%) and HESS G21.5-0.9~($\\approx$~0.02\\%).~However, all these PWNe have significantly lower levels~(ranging from 0.1 to 2.9\\%) of X-ray luminosity production efficiency than the Crab Nebula ($\\approx$~4\\%).~Despite being worst synchrotron emitters than the Crab Nebula, G54.1+0.3 together with the other recent compact TeV PWNe, are similar and even better TeV emitters than the Crab Nebula. \\par Another diagnostic is the strength of the synchrotron emission relative to the gamma-ray emission, L$_x$/L$_\\gamma$.~The Crab Nebula has the largest ratio among all young systems by far with L$_X$/L$_\\gamma$ of 220.~That of G54.1+0.3 is the lowest among the others that range between 1.4 and 30.~These numbers are consistent with~\\citet{mattana09} who finds  L$_\\gamma$$>$L$_X$ after 5 kyr from pulsar birth,~including a different list of objects and spanning over a larger range of $\\tau_c$.~The scattering of this ratio in two orders of magnitude, may be attributed to the unique environments in each of these objects, such as the local ambient photon fields or the interactions with the surrounding mediums resulting in an increase of the nebular magnetic field.~As suggested by~\\cite{chevalier04, djannati07}, these PWNe may not be in energy equipartition between particles and magnetic fields, as is the case for the Crab Nebula.~Instead, these may be particle-dominated objects.~However, proper modeling of each of these PWNe and its surroundings is needed to constrain the nebular magnetic fields. \\par Given the association by~\\citet{leahy08} of the PWN with a molecular CO cloud,~an alternative interpretation of the TeV emission from G54.1+0.3 would be radiation from the decay of neutral pions produced in the hadronic interactions produced between the dense target and the SNR. To establish if this is a compatible scenario, deeper observations of this source would be needed to constrain the size and location of the accelerated particles in the TeV range.~A detection by Fermi-LAT would add data to the spectral energy distribution of the system, giving an idea of the possible role played by protons by showing (or not) the typical bump in the spectrum expected from neutral pions (e.g.~\\citet{aharonian94}).~However, based on the association with a high spin-down pulsar, a comparable size of the X-ray and radio PWNe given the instrument resolutions and comparison with similar detected objects, the PWN scenario best explains the data.~Under the assumption that the gamma-ray production in VER J1930+188 is IC emission, the recently discovered IR shell of ejecta is an additional source of seed photons from the local dust up-scatted into the TeV energies and contributing in the observed emission."
    },
    "1005/1005.5421_arXiv.txt": {
        "abstract": "The XENON100  and CRESST experiments will directly test the inelastic dark matter explanation for DAMA's 8.9$\\sigma$ anomaly.  This article discusses how predictions for direct detection experiments depend on uncertainties in quenching factor measurements, the dark matter interaction with the Standard Model and the halo velocity distribution.  When these uncertainties are accounted for, an order of magnitude variation is found in the number of expected events at CRESST and XENON100. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.0803_arXiv.txt": {
        "abstract": "{The high-mass microquasar Cyg X-1, the best-established candidate for a stellar-mass black hole in the Galaxy, has been detected in a flaring state at very high energies (VHE), $E>$ 200 GeV, by the Atmospheric Cherenkov Telescope MAGIC. The flare occurred at orbital phase $\\phi = 0.91$, where $\\phi = 1$ is the configuration with the black hole behind the companion high-mass star, when the absorption of gamma-ray photons by photon-photon annihilation with the stellar field is expected to be  highest.} {We aim to set up a model for the high-energy emission and absorption in Cyg X-1 that can explain the nature of the observed gamma-ray flare.} {We study the gamma-ray opacity due to pair creation along the whole orbit, and for different locations of the emitter. Then we consider a possible mechanism for the production of the VHE emission.} {We present detailed calculations of the gamma-ray opacity and infer from these calculations the distance from the black hole where the emitting region was located. We suggest that the flare was the result of a jet-clump interaction where the decay products of inelastic $p-p$ collisions dominate the VHE outcome.} {We are able to reproduce the spectrum of Cyg X-1 during the observed flare under reasonable assumptions. The flare may be the first event of jet-cloud interaction ever detected at such high energies.} ",
        "introduction": "Five X-ray binaries have been detected in the very high-energy region of the electromagnetic spectrum, $E \\sim$ TeV. Three of them, PSR B1259-63, LS I +61 303 and LS 5039, have been detected at different orbital phases and show variable emission. Four gamma-ray flares were detected by the AGILE satellite from the exceptional X-ray binary Cyg X-3 (Tavani et al. 2009). The Fermi Large Area Telescope (LAT) has also detected a variable high-energy source coinciding with the position of Cyg X-3 (Abdo et al. 2009). The fifth source, Cyg X-1, has been detected only once during a flare episode. This latter detection constitutes the first evidence of very high-energy gamma-ray emission produced in the surroundings of a stellar-mass black hole (BH) in our galaxy (for further discussion see Paredes 2008).  Recently, Albert et al. (2007) reported  the results from  observations of Cyg X-1 at very high energies, $E>200$ GeV, performed with the Major Atmospheric Gamma Imaging Cherenkov (MAGIC) telescope. No persistent emission was detected, but a fast transient episode was. The satellites INTEGRAL and Swift/BAT detected with some delay a related flare at hard X-rays, while only  a statistically poor detection was found in the RXTE/ASM data at soft X-rays. This wavelength-dependent behavior may suggest that different emitting regions were involved. The gamma-ray excess occurred at orbital phase $\\phi$ = 0.91. This can help to set constraints on the location of the emission region.  More recently, the flaring nature of Cyg X-1 in gamma rays has been confirmed with the AGILE satellite (Sabatini et al. 2010). This work is devoted to a study of the absorption of high-energy photons in Cyg X-1 and the implications of the resulting constraints.  The paper is organized as follows: in the next section we describe the main characteristics of the source under study. Section 3 deals with the gamma-ray opacity by pair creation in the stellar radiation field. The production mechanism of the flare emission is then examined in the context of existing models (e.g. Bosch-Ramon, Romero, \\& Paredes 2006; Romero et al. 2003). In particular, we explore the physical conditions required by the energy budget and spectrum of the flare event. In Sect. 4 we present a simple modelization for the non-thermal emission and compare our calculations with the observational results. Finally, in Sect. 5, we present  a brief discussion and the conclusions. ",
        "conclusions": "The VHE transient emission of Cyg X-1 occurred when the BH was behind the star with respect to the observer. Because of the high absorption in the flare detection energy range, the emission close to the BH is not enough to explain the observations, unless the photons travel far away from the star, initiating a spatially extended pair cascade as considered by Zdziarski et al. (2008). This requires a fine tunned magnetic field, which allows the instantaneous isotropization of the electrons, but does not overcome their IC radiative losses. A more realistic/accurate calculation of the electromagnetic cascade propagation is then desirable. Such simulations (following the electron trajectories) will be available in a future work as an application of the code developed by Pellizza et al. (2009). Previous 1D cascade simulations (Orellana et al. 2007) are  consistent with a strong absorption and steep spectrum at TeV energies. The results by Bosch-Ramon et al (2008) have shown that if the cascades are suppressed by effects of the magnetic field, the synchrotron emission of the secondary pairs peaks at lower energies ($\\sim$ GeV).  Romero, Kaufman-Bernad\\'o, \\& Mirabel (2002) have suggested that Cyg X-1 could go through occasional microblazar phases and have estimated that the luminosity in the observer RF can be up to one order of magnitude higher than the luminosity in the jet RF. Even taking this into account, a flare triggered at the base of the jet is undetectable due to absorption at phase 0.91. A remaining option could be  a very short episode with a highly increased acretion/ejection rate, but this is speculative given the lack of evidence at  lower energies supporting the hypothesis.  Under the geometry considered here (a jet perpendicular to the orbital plane, which has an inclination of $\\sim$ 30 deg), the high-energy emission should have occurred at a large distance above the compact object where the absorbing photon field is diluted. In order to quantify the radiative outcome in this scenario we have considered the interaction of relativistic particles accelerated in a narrow region of the jet with the target particles of a dense clump of the wind.  The flare timescale is related to the permanence of the clump inside the jet. For a spherical clump with a radius $R_c$ smaller than the jet radius $R_{\\rm jet}$ $\\sim 10^{12}$ cm we can make a zerolth order estimation of the time that it takes the clump to cross the jet: $t_c$. The clump velocity is the wind velocity, which at this height is simply $v_{\\infty}$: \\begin{equation} t_c\\simeq 2R_{\\rm jet}/v_{\\infty}\\sim 10^4 \\, {\\mbox s}. \\end{equation} The flaring episode had a timescale shorter than one day and a rising time of about one hour, which is on the same order as the $t_c$ estimated.  The simple model presented here for the broadband spectrum of Cygnus X-1  reproduces fairly well the observed  SED by MAGIC during the flare using a set of parameters that agrees with reasonable  values for this source. Interactions between the clumpy winds of massive stars with the relativistic jets in HMMQ are expected to be  produce flaring episodes at high and very high energies, and may be detectable by the new high-energy detectors, like Fermi, MAGIC II, and VERITAS."
    },
    "1005/1005.2127_arXiv.txt": {
        "abstract": "We provide a both qualitative and quantitative comparison among different approaches aimed to solve the problem of non-linear diffusive acceleration of particles at shocks. In particular, we show that state-of-the-art models (numerical, Monte Carlo and semi-analytical), even if based on different physical assumptions and implementations, for typical environmental parameters lead to very consistent results in terms of shock hydrodynamics, cosmic ray spectrum and also escaping flux spectrum and anisotropy. Strong points and limits of each approach are also discussed, as a function of the problem one wants to study. ",
        "introduction": "In the late 70's many authors \\citep{krimskii,als77,b078,bell78} introduced the \\textit{test-particle} theory of particle acceleration at strong collisionless shocks due to first order Fermi mechanism. However, quantitative estimates soon pointed out that this mechanism, called Diffusive Shock Acceleration (DSA), may be so efficient that the back-reaction of the accelerated particles on the shock dynamics can not be neglected.  The obvious theoretical challenge is how to model effectively the full shock dynamics. Modeling shocks with efficient particle acceleration using plasma simulations (PIC and hybrid) is extremely computationally expensive for several reasons. First, the energies of charged particles participating in the process range from the low thermal energies of cold plasma to the ultra-relativistic energies of cosmic rays, and both extremes of particle spectra are dynamically important if acceleration is efficient \\citep[see, e.g.,][for estimates of computational requirements]{veb08}. Second, simulations need to be done in three dimensions because of the possibility of non-physical suppression of important processes in one- and two-dimensional simulations \\citep[see, e.g., ][]{jjb98}. Therefore, approximate methods must be used to model efficient particle accelerating shocks.  The first step towards a Non-Linear theory of DSA (NLDSA), in which particle acceleration and shock dynamics are calculated self-consistently, is represented by the two-fluid model proposed by \\cite{dr_v80,dr_v81}. This approach is indeed very useful to study the dynamics of a shock modified by the presence of accelerated particles (treated as a diffusive, relativistic fluid) but by design it cannot provide any information about the spectrum of the cosmic rays (CRs) which are produced.  A \\textit{kinetic} study of the NLDSA at astrophysical shocks also requires a statistical description of transport of superthermal particles. This relevant piece of Physics has been taken into account by adopting a number of rather different approaches; namely \\begin{itemize} \\item fully numerical simulations, in which a time-dependent diffusion-convection equation for the CR transport is solved in concert with coupled gas dynamics conservation laws, like in \\cite{bell87, bv97, kj97, kjg02, kj05, kj06, kj07, za10}; \\item stationary Monte Carlo numerical simulations of the full particle population, like in \\cite{ee84,ee85,emp90,ebj95,ebj96,ed02,veb06}; \\item semi-analytic solutions of the stationary \\citep[or quasi-stationary, see][]{bac07} diffusion-convection equation coupled to the gas dynamic equations, like in \\cite{malkov97,mdv00, blasi1, blasi2, ab05, ab06,cabv09,cab09}. \\end{itemize}  All these approaches make consistent predictions about the main consequences of the shock modification, i.e.\\ the appearance of an upstream precursor created  by the CR pressure around the shock, which slows down the incoming fluid. As a consequence of this CR-induced velocity gradient, the fluid compression ratio is no longer 4 for any strong shock. Hence, the test-particle prediction that the spectrum of accelerated particles has to be $\\propto p^{-4}$ is no longer valid. Nevertheless, it is still true that the spectral slope mainly depends on the compression factor actually ``felt'' by CRs. Since the larger is its momentum, the farther from the shock a particle can diffuse, high (low) momentum particles probe a total compression ratio larger (smaller) than 4. The resulting CR spectrum becomes rather concave, being harder (softer) than $p^{-4}$ at the higher (lower) energies. For accurate reviews of CR-modified shocks the reader can refer to \\cite{drury83,be87,je91,md01}.  In this paper we want, for the first time, to construct both  qualitative and quantitative comparisons among different approaches to NLDSA, in  order to highlight their analogies and differences. In particular, we compare three state-of-the-art kinetic methods which are representative of the three categories listed above: the fully numerical calculation by \\cite{kj07}, the Monte Carlo simulation by \\cite{veb06} and the semi-analytical solution by \\cite{cab09}. The final goal of this work is to provide a useful tool illustrating which physical aspects of the problem can be taken into account within each approach, which assumptions are \\textit{a posteriori} justified and consistent with any observational signature, and finally what is the computational effort required to tackle a given problem.  The present work is organized as follows: the next three sections are dedicated to the description of the three different techniques adopted to solve the problem of NLDSA. For each of them, the assumptions (e.g.\\ time-dependence, isotropy of the CR distribution function) and the equations of the model, the injection of particles in the acceleration mechanism and the turbulent heating in the precursor are outlined. In \\S\\ref{sec:benchmark} the results obtained within different approaches are compared, in terms of both hydrodynamics and CR distribution function. We comment and conclude in \\S\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} In this work we conducted a comparison among numerical, Monte Carlo and semi-analytical approaches to NLDSA theory by solving a benchmark case for an astrophysical plane, non-relativistic strong shock. In particular we examined a case in which cosmic rays are efficiently accelerated and strongly modify the shock with their dynamical back-reaction.  The results obtained with the three different techniques that assume specific and distinct recipes for physical ingredients such as particle injection, particle escape and diffusion properties are very consistent with each other in terms of both accelerated (and escaping) spectra and hydrodynamics. These findings represent an actual investigation of the role that specific pieces of information embedded in each approach have in modeling cosmic ray modified shocks. In particular, we showed how the solution obtained when the stationary state is achieved by (a) letting the system evolve at the presence of a free-escape boundary (the numerical approach, see\\S\\ref{sec:KJ}), (b) a steady-state solution in which the particle transport is described statistically through a Monte Carlo technique (\\S\\ref{sec:EV}), or (c) the semi-analytic solution of the stationary diffusion-convection equation for the isotropic part of the cosmic ray distribution function (\\S\\ref{sec:CAB}) show an agreement well within any uncertainty intrinsic to the basic physics of the process. Thus, the agreement would not allow distinction of the methods using current observations.  At the same time, it is worth recalling that all the models are rather flexible and may be updated with any new physical insight which may come from present and future observations. Very generally, these methods are limited more by our real knowledge of the mechanisms going on in a SNR shock than by the possibility of implementation of these effects.  An important example of a piece of information that can be embedded in all the methods is a more detailed description of the magnetic field responsible for the particle scattering. In the calculations above we did not include any process leading to the amplification of  the background magnetic field due to CR streaming instabilities \\citep{bell78, bell04}, nor any dynamical feedback of the magnetic turbulence in the shock dynamics \\citep{cbavlett}. On the other hand, by assuming Bohm-like diffusion we have intrinsically taken into account a magnetic configuration quite different from the typical interstellar one \\citep{lc83a,lc83b}. Moreover, we included a phenomenological model for a non-adiabatic heating of the upstream background plasma as due to the instantaneous damping of all the magnetic turbulence produced via resonant streaming instability, as described in \\cite{mck-v82}. This very efficient damping, while largely adopted in order to prevent dramatic shock modification, is however inconsistent with magnetic fields of order $100-500\\mu$G inferred in the downstream of some young SNRs \\citep[see e.g.][]{P+06}: in these environments the shock modification may be controlled by the non-negligible dynamical role of the amplified magnetic field, an effect which can be easily included in the models above \\citep[see][]{veb06,cabv09}.  Another interesting finding of our comparison is that not only Monte Carlo simulations, but also methods based on the solution of the diffusion-convection equation can retain reliable information about the anisotropy of the distribution function, even close to the free-escape boundary, in many astrophysical cases. This fact allows one to use the so-derived diffusive CR current to calculate the excitation of modes relevant for the growth of the magnetic turbulence, as predicted in many scenarios involving the self-generation of an amplified magnetic field responsible in which particles are scattered very efficiently \\citep[see e.g.][for the calculation of the Bell non-resonant streaming instability excited by escaping particles]{reville}.  In this work we considered a situation in which injection occurs very efficiently, as suggested by the Monte Carlo modeling of the subshock. This choice, while on one hand perfectly reproduces the observations of the Earth bow shock \\citep[e.g.][]{emp90}, on the other hand is not firmly supported by observational evidence in SNR shocks. For typical values of temperatures and magnetic fields in SNR environments, in fact, the region of the spectrum between the thermal and the non-thermal distributions does not provide, at the moment, any strong signature testable with multi-wavelength studies of SNRs. In this respect, a simple parametrization of the injection process like the thermal leakage captures most of the (non-linear) physics which is expected to be relevant in NLDSA theory. Further improvements in discrimination between the various models of particle injection can be expected from the ongoing work on particle-in-cell simulations of non-relativistic shocks.  Dealing with strongly modified shocks is indeed very challenging for fully numerical methods. When huge gradients are found in the precursor, the system very likely develops some instabilities, whose nature may be not only numerical, but also physical.  For instance, the acoustic instability studied by \\cite{df86,kjr92,w+09} can develop strong inhomogeneities, which can drive turbulence in multi-dimensional flows. Numerical methods are thus tools necessary in order to unravel these kind of questions, being in principle able to follow the growth of instabilities even in the non-linear regime, typically inaccessible to analytical approaches.  A numerical simulation is currently the only way to deal with time-dependent NLDSA problems. We recall that there are two main ways of accounting for the dependence on time of the physical quantities: on one hand we have cases in which the shock velocity is almost constant and only CR-related quantities are expected to vary on time scales comparable with the acceleration time at $p_{max}$, and on the other hand we have cases in which also ``environmental'' quantities, like the shock velocity, change with time.  The former case is realized for instance during the ejecta dominated stage of a SNR, in which the shock is not dramatically slowed down by the inertia of the swept-up material. Nevertheless, if a spatial boundary condition is imposed in such a way that particles are free to escape from an upstream boundary, also in this case a steady-state solution consistent with stationary approaches may be recovered. Otherwise, if $p_{max}$ is free to grow unconstrained by spatial limits, the system evolves according to a self-similar behavior as shown by \\cite{kj07,krj09}: also in this case stationary approaches, with a proper age-limited maximum momentum \\citep[see]{bac07}, should be able to reproduce the correct results.  The latter case, instead, is expected to be realized in the Sedov-Taylor stage of the evolution of SNRs. As shown by \\cite{escape}, during this stage the decrease of the shock velocity is expected to lead to a decrease of the magnetic field amplification and in turn to a drop of the confinement power of the system. At any given time particles with momentum close to $p_{max}(t)$ cannot diffuse back to the shock and escape the system upstream, thus, carrying away a sizable fraction of the energy flow into the shock. This situation is thought to lead to the achievement of a quasi-stationary configuration in which particle acceleration and losses balance themselves; also in this case a time-independent approach might be a good description of the problem.  In the case in which the assumption of quasi-stationarity is satisfied, modeling an astrophysical object might require one to run the NLDSA code many times: one for each time-step, for any given choice of the many environmental and model parameters. It is clear that such an investigation can be conducted only by using very fast solutions of the NLDSA problem, and semi-analytical approaches can do this job very well \\citep[see e.g.][for an example]{cab10}.  Finally, Monte Carlo codes allow one to take into account strong anisotropies of the CR distribution function, for which the diffusive approximation and hence the solution of the diffusion-convection equation for the isotropic part of $F(x,{\\bf p})$ only may fail by a large amount. This point is of central interest especially for the study of mildly-relativistic and relativistic shocks \\citep{ed02}, since anisotropic corrections to Eq.~\\ref{diffcon} are of order $\\mathcal{O}(u^{2}/c^{2})$, as outlined in \\S\\ref{sec:escape}.  The last consideration is about an assumption that is ubiquitous in NLDSA theories; namely, that particle are efficiently scattered by magnetic irregularities and in a \\textit{diffusive} way. This assumption may be broken close to $x_{0}$, beyond which particles are expected to stream away without diffusing, or even in the precursor if the topology of the magnetic field is organized in structures like the cavities predicted by Bell's instability \\citep{bell04,bell05}. Nonetheless it is worth stressing that in the literature no non-linear particle acceleration methods able to relax the hypothesis that particle are scattered in a diffusive way (which is embedded also in Monte Carlo approaches) have been developed yet. Such a problem could be in principle addressed by Particle in Cell (PIC) simulations, but the computational time required to run a case in which the spectrum of accelerated particles spans many order of magnitudes is now, and it is going to be for years, prohibitive even for huge computer clusters."
    },
    "1005/1005.0122_arXiv.txt": {
        "abstract": "We present the Hubble diagram (HD) of 66 Gamma Ray Bursts (GRBs) derived using only data from their X\\,-\\,ray afterglow lightcurve. To this end, we use the recently updated $L_X$\\,-\\,$T_a$ correlation between the break time $T_a$ and the X\\,-\\,ray luminosity $L_X$ measured at $T_a$ calibrated from a sample of {\\it Swift} GRBs with lightcurves well fitted by the Willingale et al. (2007) model. We then investigate the use of this HD to constrain cosmological parameters when used alone or in combination with other data showing that the use of GRBs leads to constraints in agreement with previous results in literature. We finally argue that a larger sample of high luminosity GRBs can provide a valuable information in the search for the correct cosmological model. ",
        "introduction": "That the universe is spatially flat, has a subcritical matter content and is undergoing a phase of accelerated expansion are nowadays globally accepted ideas. Strong evidence for this scenario come from the anisotropy and polarization spectra of the Cosmic Microwave Background Radiation (CMBR) \\cite{Boom,QUAD,WMAP7}, the galaxy power spectrum \\cite{Teg06,P07} with its Baryonic Acoustic Oscillations (BAO) \\cite{Eis05,P10} and the Hubble diagram (HD) of Type Ia Supernovae (SNeIa) \\cite{Union,H09,SNeIaSDSS}. What is driving this cosmic speed up and dominating the energy budget is, on the contrary, a still hotly debated question with answers running from the classical cosmological constant \\cite{CPT92,SS00}, scalar fields \\cite{PR03,Copeland} and higher order gravity theories \\cite{CF08,FS08,ND08,dFT10}. In order to break the degeneracies among model parameters and go to the next level of model selection, it is mandatory to probe the background evolution up to very high redshift so that the onset of the acceleration epoch and the transition to the matter dominated regimes may be directly investigated.  \\begin{figure*} \\centering \\subfigure{\\includegraphics[width=7cm]{avsom.eps}} \\goodgap \\subfigure{\\includegraphics[width=7cm]{bvsom.eps}} \\goodgap \\\\ \\subfigure{\\includegraphics[width=7cm]{svsom.eps}} \\goodgap \\subfigure{\\includegraphics[width=7cm]{rmsvsom.eps}} \\goodgap \\\\ \\caption{Best fit calibration parameters $(a, b, \\sigma_{int})$ and rms of the residuals for the $L_X$\\,-\\,$T_a$ correlation fitted to the 8 canonical GRBs as a function of the matter density parameter $\\Omega_M$ for three different values of $w_0$, namely $w_0 = -1.25$ (short dashed), $w_0 = -1.0$ (solid), $w_0 = -0.75$ (long dashed). We use a CPL model with $h = 0.70$ and set $w_a = -w_0$ to have a matter dominated era at high redshift.} \\label{fig: cosmopar} \\end{figure*}  Thanks to their enormous energy release, GRBs are visible up to very high $z$, the largest one being at $z = 8.2$ \\cite{Salvaterra2009}, hence appearing as ideal candidates for this task. Unfortunately, GRBs are everything but standard candles because their peak luminosity spans a wide range. There have nevertheless been many attempts to make them standardizeable candles resorting to the use of empirical correlations among distance dependent quantities and rest frame observables \\cite{Amati08,FRR00,N000,G04,liza05}. Such empirical relations allow one to infer the GRB rest frame luminosity or energy from an observer frame measured quantity so that the distance modulus can be estimated with an error mainly depending on the intrinsic scatter of the adopted correlation. Combining the estimates from different correlations, Schaefer (2007) first derived the GRBs HD for 69 objects, while Cardone et al. (2009) used an enlarged sample and a different calibration method to update the GRBs HD. Many attempts on using GRBs as cosmological tools have since then been performed (see, e.g., Firmani et al. 2006, Liang et al. 2009, Qi \\& Lu 2009, Izzo et al. 2009 and refs. therein) showing the interest in this application of GRBs.  In this {\\it Letter}, we rely on the $L_X$\\,-\\,$T_a$ correlation \\cite{D08} to build the HD of 66 GRBs from their X\\,-\\,ray afterglow lightcurves observed with the {\\it Swift} satellite \\cite{Evans2009}. We then present a preliminary application of the derived GRBs HD showing how this dataset can help constrain the parameters of some simple dark energy models. ",
        "conclusions": "A {\\it ground zero} test of every cosmological model is its ability to fit the observed Hubble diagram. Such a test tells us whether the model is able to give rise to the observed background evolution and select the region of the parameter space leading to the correct sequence of accelerating and decelerating expansion. Although still being the primary tracer of the Hubble diagram, SNeIa may trace cosmic evolution only up to $z \\sim 2$ so that it is mandatory to look for a different class of astronomical objects to go beyond this limit and investigate the matter dominated era. GRBs are considered the ideal candidates for this role thus motivating the hunt for a calibration method to make them standardizeable candles. We have presented here the use of the $L_X$\\,-\\,$T_a$ correlation as a valid method to infer the GRBs HD up to $z = 8.2$ using data coming from the X\\,-\\,ray lightcurve. It is worth stressing that this is the only empirical law relating quantities measured from the afterglow lightcurve rather than being related to the prompt emission quantities. Moreover, differently from what has yet been done in the past \\cite{S07,CCD09}, the HD for the NC\\,+\\,C sample is the only GRBs HD based on the use of a single correlation and containing a statistically meaningful number of objects. The use of the $L_X$\\,-\\,$T_a$ correlation then avoids the need of combining different correlations to increase the number of GRBs with a known distance modulus. Each correlation is affected by its own possible systematics and characterized by different intrinsic scatter so that combining all of them in a single HD can introduce unexpected features and hence bias the constraints on the cosmological parameters.  The intrinsic scatter of the $L_X$\\,-\\,$T_a$ correlation may be significantly reduced if one considers only its upper envelope defined by the canonical GRBs. On the other hand, using the full 66 GRBs, we have been able to get constraints on the matter content and the present day Hubble constant comparable to those yet available in literature. Such a preliminary investigation has convincingly shown that the $L_X$\\,-\\,$T_a$ correlation may be used to construct a GRBs HD which does not introduce any bias in the search for cosmological parameters. It is worth stressing that the same results have been obtained using the canonical GRBs despite the fact that they represent just $\\sim 12\\%$ of the full sample. We therefore argue that an observational effort dedicated to look for canonical GRBs may turn them from ideal candidates to actual tools for efficiently investigating the dark energy puzzle. \\\\  {\\it Acknowledgements.} We warmly thank M. Ostrowski for his valuable comments. This work made use of data supplied by the UK Swift Science Data Centre at the University of Leicester. MGD is grateful for the support from Polish MNiSW through the grant N N203 380336. MGD is also grateful for the support from Angelo Della Riccia Foundation."
    },
    "1005/1005.1312_arXiv.txt": {
        "abstract": "Both the acceleration of cosmic rays (CR) in supernova remnant shocks and their subsequent propagation through the random magnetic field of the Galaxy deem to result in an almost isotropic CR spectrum. Yet the MILAGRO TeV observatory discovered a sharp ($\\sim10^{\\circ})$ arrival anisotropy of CR nuclei. We suggest a mechanism for producing a weak and narrow CR beam which operates en route to the observer. The key assumption is that CRs are scattered by a strongly anisotropic Alfven wave spectrum formed by the turbulent cascade across the local field direction. The strongest pitch-angle scattering occurs for particles moving almost precisely along the field line. Partly because this direction is also the direction of minimum of the large scale CR angular distribution, the enhanced scattering results in a weak but narrow particle excess. The width, the fractional excess and the maximum momentum of the beam are calculated from a systematic transport theory depending on a single scale $l$ which can be associated with the longest Alfven wave, efficiently scattering the beam. The best match to all the three characteristics of the beam is achieved at $l\\sim1$pc. The distance to a possible source of the beam is estimated to be within a few 100pc. Possible approaches to determination of the scale $l$ from the characteristics of the source are discussed. Alternative scenarios of drawing the beam from the galactic CR background are considered. The beam related large scale anisotropic CR component is found to be energy independent which is also consistent with the observations. ",
        "introduction": "}  The MILAGRO TeV observatory recently discovered collimated beams dominated by hadronic cosmic rays (CR) with a narrow ($\\sim10^{\\circ}$) angular distribution in the 10 TeV energy range \\citep{Milagro08PRL}. This is surprising, since most of the CR acceleration and propagation models predict only a weak, large scale anisotropy. The acceleration models are based on the diffusive shock acceleration (DSA) mechanism, widely believed to generate galactic CRs in supernova remnant shocks (SNR). The corner stone of the DSA is a rapid pitch-angle scattering of CRs by self-generated Alfven waves in the shock vicinity. An enhanced scattering isotropizes particle distributions. Moreover, when the shock releases the accelerated particles into the interstellar medium (ISM), they continue to scatter by the ISM turbulence. Even though this scattering occurs at a significantly lower rate, all sharp anisotropies carried over from the accelerator or created otherwise, should be erased during the long ($\\ga$100 pc) travel of the CRs from any hypothetical nearby SNR to the observer. Yet, the astounding sharp beaming effect is argued to be genuine.  Focusing on relatively distant accelerators (such as nearby SNRs) and long-distance propagation effects as a possible cause of the MILAGRO beam(s), we do not consider 'local' scenarios that have already been discussed and largely rejected by \\citet{DruryAharMILAGRO08} and \\citet{SalvatiMilagro08}. As for the remote accelerator with subsequent propagation effects, some of them have also been suggested in the above publications. In particular, \\citet{SalvatiMilagro08} associate the observed CR beam with the Geminga pulsar. However, \\citet{DruryAharMILAGRO08} argue that this does not explain namely the sharp collimation, and suggest a magnetic nozzle as such a collimation device. The magnetic nozzle scenario, however, poses a rather strong constraint on the nozzle mirror ratio ($B_{max}/B_{min}\\sim\\vartheta^{-2}\\gg1,$ where $\\vartheta$ is the beam angular width). The advantage of this scenario is that the beam density is equal to the difference between the isotropic components of CRs on each side of the mirror (by linearity of the transport equation). Since the Milagro beam is very weak ($\\sim10^{-4}$ of the CR density), this is a very mild requirement on the CR enhancement on the far side of the mirror. It is also true that the existence of a magnetic mirror of that strength cannot be warranted or denied on rational ground. It should be noted that any anisotropic distribution may become vulnerable to self-spreading in pitch angle. As pointed out by \\citet{DruryAharMILAGRO08}, however, the isotropic CR background should produce a stabilizing effect against the beam self-spreading. We will quantify the CR stabilization in Sec.\\ref{sec:BeamSustainability} required for both the collimation mechanism suggested in the present paper and for the magnetic nozzle hypothesis.  In this paper we suggest a novel mechanism for producing a narrow CR beam. It is based on the strong anisotropy of the MHD turbulence in the ISM. Such anisotropy is expected when the turbulence is driven at a long (outer) scale, but unlike the isotropic Kolmogorov cascade, the incompressible MHD cascade is directed perpendicularly to the magnetic field in the wave vector space. This was shown by \\citet{Goldr95} (GS) (see also \\citealt{SridhGoldr94} and \\citealt{goldr97}) and confirmed by numerical simulations (\\eg \\citealt{ChoVishn00,MaronGoldr01,BereznLaz09}). The cascade proceeds to $k_{\\perp}r_{g}\\left(p\\right)\\gg1$ in the perpendicular wave number direction for the protons with the gyro-radii $r_{g}\\sim10^{16}$cm, typical for the particles of the MILAGRO beam energies $pc\\sim10$TeV and the ISM magnetic field of a few $\\mu G$. Contrary to the $k_{\\perp}$ direction the spectrum spreading in $k_{\\parallel}$ is suppressed, so that $k_{\\parallel}\\sim k_{\\perp}^{2/3}l^{-1/3}\\ll k_{\\perp}$, where $l$ is the outer scale.  As is known from the wave-particle interaction in plasmas, the scattering of particles with the Larmor radius exceeding the wave length in the perpendicular direction, $k_{\\perp}r_{g}\\gg1$ is strongly suppressed, since such particles suffer a rapidly changing electromagnetic force. Specifically, the CR scattering by the GS anisotropic spectrum was investigated in a number of publications (\\citealt{ChandranGS00PhRvL,YanLazar02}). What is important for the purposes of this paper is that the pitch-angle scattering rate is peaked at $\\left|\\mu\\right|=\\left|\\cos\\vartheta\\right|\\approx1$, \\ie for particles moving along the field line, since for these particles $k_{\\perp}r_{g}\\left(p_{\\perp}\\right)\\la1$. Only particles with such small $p_{\\perp}$, \\ie with pitch angles within $\\sin^{2}\\vartheta\\la\\epsilon\\ll1$ are scattered efficiently. Looking at this problem mathematically, a peaked diffusion coefficient $D\\left(\\mu\\right)$ does not necessarily result in a peaked particle distribution $f\\left(\\mu\\right)$. Indeed, the time-asymptotic solution of the diffusion equation with zero flux through the boundary is a flat distribution even if the diffusion coefficient is not constant. Nevertheless, consider the particle diffusion in pitch angle on an\\emph{ intermediate time-scale}, \\ie  when anisotropy is erased within the strong peak of the diffusion coefficient $D\\left(\\mu\\right)$, but is present in the region where $D\\left(\\mu\\right)$ is much smaller. The dominant eigenfunction of the scattering operator has a relatively broad minimum at $\\mu=1$, \\ie where $D$ is sharply peaked. Now, the enhanced scattering fills up the very bottom of this minimum which appears as a narrow excess, Fig.\\ref{fig:Pitch-angle-dependence}. In the context of a classical Lorentz gas relaxation problem \\citep{Gurevich61RunAway,KruskalBernstein64}, this is clearly a transient effect associated with an incomplete decay of the anisotropic part of the pitch-angle distribution. Note that the difference with the Lorenz gas problem is in the sharply peaked $D\\left(\\mu\\right)$. In addition, our problem is a problem in $z$, which is the spatial coordinate (rather than time) along a magnetic flux tube that connects the CR source with the Earth.  The demonstration of this phenomenon, facilitated and obscured at the same time by the fact that the peak region $\\left|\\sin\\vartheta\\right|\\ll1$ contains the singular points of particle pitch-angle diffusion operator at $\\sin\\vartheta=0$, will be the main subject of the present paper.  Before we tackle this problem, we briefly discuss in the next section how narrow the CR angular distribution can be as it leaks from a hypothetical nearby SNR accelerator, magnetically connected with the heliosphere. This, or some other moderately anisotropic distribution of CRs, created by a recent SNR explosion, will be subjected in Sec.\\ref{sec:Pitch-angle-scattering-of} to the pitch-angle scattering analysis and to the propagation analysis in Sec.\\ref{sec:Particle-propagation}. Next, in Sec.\\ref{sec:BeamSustainability} we determine the maximum energy of the beam beyond which it must spread on self-generated Alfven waves. Sec.\\ref{Sec:Distance-to-the} deals with the relation between the beam maximum energy and the distance to its possible source. We conclude with a brief discussion of the results and of what the fascinating MILAGRO findings can possibly tell about a nearby accelerator and the structure of ISM turbulence. ",
        "conclusions": "}  The principal results of this paper are as follows. Assuming only a \\emph{large scale} anisotropic distribution of CRs (generated, for example by a nearby accelerator, such as a SNR) and a \\citet{Goldr95} (GS) cascade of Alfvenic turbulence originating from some scale $l$, which is the longest scale relevant for the wave-particle interactions, we calculated the propagation of the CRs down their gradient along the interstellar magnetic field. It is found that the CR distribution develops a characteristic angular shape consisting of a large scale anisotropic part (first eigenfunction of the pitch-angle scattering operator) superposed by a beam, tightly focused in the momentum space in the local field direction. The large scale anisotropy carries the \\emph{momentum dependence }of the source,\\emph{ }while both the beam angular width and its fractional excess (with respect to the large scale anisotropic component) grow with momentum (as $\\sqrt{p}$ and $p$, respectively). Apart from the width and the fractional excess of the beam, the theory predicts its maximum momentum on the ground that beyond this momentum the beam destroys itself. All the three quantities are completely determined by the turbulence scale $l$. Even if $l$ is considered unknown, it can be inferred from any of the three independent MILAGRO measurements. These are the width, the fractional excess and the maximum energy of the beam, and all the three consistently imply the same scale $l\\sim$1 pc. The calculated beam maximum momentum encouragingly agrees with that measured by MILAGRO (\\textasciitilde{}10 TeV/c). The theoretical value for the angular width of the beam is found to be $\\Delta\\vartheta\\simeq4\\sqrt{\\epsilon}$, where $\\epsilon=r_{g}\\left(p\\right)/l\\ll1$. The beam fractional excess related to the large scale anisotropic part of the CR distribution is $\\simeq50\\epsilon$. Both quantities also match the Milagro results for $E\\sim1-2$ TeV. So the beam has a momentum scaling that is one power shallower than the CR carrier, it is drawn from. This finding will receive a due discussion.  Obviously, the determination of the absolute value of the beam excess would require the source intensity. For the lack of such information, an indirect inference was made in Sec.\\ref{sub:Observational-consequences} about the galactic CR (GCR) large scale anisotropy being of the same order as the large scale anisotropy responsible for the beam.  Below the rationale for this admission is given which we open with the following notations: \\begin{itemize} \\item $\\tilde{F}_{GCR}$ and $\\bar{F}_{GCR}$ are the large scale anisotropic and isotropic parts of the galactic (not assumed to be related to the source of the beam) distribution, respectively \\item $\\tilde{F}_{S}$ and $\\bar{F}_{S}$ (i.e. $f_{\\infty}$ in sec.\\ref{sec:Particle-propagation}) are the similar quantities related to the source of the beam at the distance $L_{S}$ \\item $F_{B}$ is the beam distribution on the top of $\\tilde{F}_{S}$ \\end{itemize} Unless the source of the beam is also responsible for the GCR, the quantities $\\tilde{F}_{S}$ and $\\tilde{F}_{GCR}$ are independent of each other and cannot be related since the source intensity, the distance to it, $L_{S}$ and the losses are unknown. Even if the beam and its carrier $\\tilde{F}_{S}$ propagate without significant losses (or suffer similar losses), the current theory determines only a fractional excess $F_{B}/\\tilde{F}_{S}\\sim50\\epsilon$ (independent of $L_{S}$). For the same reasons we do not know what is the source contribution $\\bar{F}_{S}$ into the total isotropic CR background $\\bar{F}{}_{GCR}+\\bar{F}_{S}$. However, since the beam A is not observed at the minimum% \\footnote{It is interesting to note that \\citet{Milagro08PRL} point out that there is a deep deficit bordering the excess regions. This deficit could be identified as a minimum of the dominant eigenfunction, but they attribute it to the effect of including the excess regions into the background. In other words the deficit is an artefact of the data analysis.% } of the (measured) total large scale $\\tilde{F}_{S}+\\tilde{F}_{GCR}$ as it would, were $\\tilde{F}_{S}\\gg\\tilde{F}_{GCR}$ the case (see Fig.\\ref{fig:Eigenfunction}), we infer $\\tilde{F}_{GCR}\\ga\\tilde{F}_{S}$. Furthermore, since $F_{B}/\\tilde{F}_{S}$ is calculated and $\\tilde{F}_{S}+\\tilde{F}_{GCR}$ is measured along with $F_{B}$, the both quantities $\\tilde{F}_{S}$ and $\\tilde{F}_{GCR}$ can also be determined.  We found that $F_{B}/\\tilde{F}_{S}\\sim0.1$ for $l\\sim1$pc which was, in turn, deduced from two other independent measurements (beam width $\\Delta\\vartheta$ and its maximum energy $E_{Bmax}$). Since Milagro measurements indicate that $F_{B}/\\left(\\tilde{F}_{GCR}+\\tilde{F}_{S}\\right)\\sim0.1$, we conclude that $\\tilde{F}_{S}\\sim\\tilde{F}_{GCR}$. A more specific relation between the two would not be meaningful since the measurements of $F_{B}\\left(p\\right)$ are rather limited. If particle losses from the flux tube are negligible, it follows that $\\tilde{F}_{S}\\left(L_{S}\\right)\\sim\\tilde{F}_{S}\\left(0\\right)$ for $p\\ll p_{Bmax}$ ($0$- being the source position).  These findings allow us to speculate about the possible source of the beam. First, if the source is an active accelerator that emits strongly anisotropic particle flux, the last relation implies that $\\bar{F}_{S}\\sim\\tilde{F}_{S}$. Since by observations $\\tilde{F}_{S}\\ll\\bar{F}_{GCR}$, such source cannot contribute significantly to the 'knee' region at $\\simeq3$ PeV. In this case our inferences of $l$ from three independent measurements --all strikingly pointing at the $3$ PeV accelerator cut-off energy (with $l\\sim r_{g}\\left(E_{max}\\right)$)-- must be either a coincidence or a different mechanism couples the galactic 'knee' particles with the scale of the turbulence that generates the beam. If it is not a coincidence and the source contributes significantly to the observed CR background, the escaping particle flux should be quasi-isotropic, $\\tilde{F}_{S}\\ll\\bar{F}_{S}$ (to allow for $\\bar{F}_{S}\\sim\\bar{F}_{GCR}$). In combination with the assumption that particles escape in the wide range 1TeV -3 PeV (to both form the beam and to inject MHD energy at the scale $l$), the source is unlikely to be an active accelerator, but rather a region of an enhanced CR density, with a steep cut-off at $\\simeq3$ PeV. The near isotropy at the source is not inconsistent with a currently working accelerator, but escape in such a broad energy range probably is. Indeed, at least the available (known to us) mechanisms, that offer a broad energy escape from a SNR along with the spectrum steepening (\\ie spectral break, starting 1-2 orders of magnitude below the cut-off, \\eg \\citealp{MDS05,MD06}) seem to fall short to cover three orders of magnitude in energy. Moreover for the source to be a recent accelerator (such as a recent SNR, suggested by \\citealp{ErlykinWolfendale97} with the spectrum $E^{-2}$) the mechanism should be found that makes the spectrum of the escaping particles at least $0.5$ steeper (and still steeper if the acceleration was strongly nonlinear). Combined with the nonlinear acceleration (which is required in \\citealp{MD06} model) this would make an acceptable spectrum but again, it is not clear how these particles can initiate the MHD cascade at such a long scale, to ensure the required value of $l$.  Our argument against the beam and the bulk CR $\\tilde{F}_{GCR}$ coming from the same source is that the observed beam is not located at the minimum of the angular distribution of the first eigenfunction, so we need to allow for a second component. This is largely a technical limitation, stemming from 1-D transport model, in which $F_{B}$ and $\\tilde{F}_{S}$ are coupled, as wells as from the single eigenfunction approximation. By removing this latter simplification alone (which is probably even necessary for an accurate description of $\\tilde{F}_{S}$ at TeV- energies, Sec.\\ref{Sec:Distance-to-the}) the above constraint can be relaxed. Another possibility is a lateral diffusion and drifts of $\\tilde{F}_{S}$-component from the flux tube.  An interesting obvious conjecture from the common origin of the beam and $\\tilde{F}_{GCR}$ would be that the proton 'knee' at $\\simeq3$ PeV is also of the same origin as the beam. However, the beam spectrum is calculated to be one power flatter than its carrier. According to Milagro the beam index is about $1.5$, so that the carrier should have an index $\\simeq2.5$ which is closer to the GCR than to a hypothetical 'recent SNR'. In particular, this would not support the single source hypothesis of the GCR 'knee' \\citet{ErlykinWolfendale97}. Equally problematic would be an active accelerator scenario, unless the steepening mechanisms of the run-away CR mentioned above can be adopted after due modifications.  All told, the beam is likely to be at least partly drawn from the GCR (due to the relation between the indices $q_{B}\\simeq q_{GCR}-1$) but the GS- turbulence that creates the beam must be driven by considerably more energetic, $\\ga1$ PeV particles (due to the constraint, $l\\simeq r_{g}$), unless the spiral-arm 1pc value \\citet{OuterScale08} for $l$ is, indeed, acceptable. To explore the possibility of the GCR origin of the beam, an extension of the above model is necessary. At a minimum, the model should include the transport of energetic particles across the flux tube. On the other hand, the particle beaming processes should remain similar to that described in Sec.\\ref{sec:Pitch-angle-scattering-of}. However, such consideration is out of the scope of this paper, particularly because the transport across the flux tube requires a separate study.  Yet another possibility is that the required GS cascade starts in the local interstellar cloud (LIC). It has a suitable size of $\\sim5$ pc \\citep{LIC00} and there would be no problem with the spectrum slope since the beam would be drawn from the GCR with the 'right' spectral index $q_{GCR}\\simeq2.7$. Whether the turbulence energy can be injected at the required scale remains to be studied. If it can, the above transport and beam focusing mechanism would be applicable since a parsec wave length and the GS cascade are the only requirements to draw the beam out of the background CR distribution.  To conclude, the model presented in this paper offers an explanation of the most pronounced Milagro beam A, while there are two more. One of them is the beam B, $\\sim50^{\\circ}$ away from beam A and the second one is in the Cygnus loop area $\\sim100^{\\circ}$ away. Any attempt to incorporate those two beams into our current model would be speculative. We merely note that the local ISM environment is complicated indeed thus offering many possibilities in explaining various CR anomalies (\\eg \\citealp{Amenomori07}). Approaches to the explanation of all three beams based on such a complexity, could hardly pass the Occam's razor test. In contrast, the model suggested in this paper is devoid of free parameters, if the knee energy at $\\sim3$PeV can be associated with the maximum CR energy of the source of the Beam A. Even though such an association is not proven, our propagation model predicts the following three beam characteristics: its width, fractional excess and maximum energy to be the functions of a single quantity, the longest wave-particle interaction scale $l$. They all give the correct MILAGRO values for $l\\simeq1$ pc, which is unlikely to be coincidental. However, the exact origin of this particular value remains unclear."
    },
    "1005/1005.1638_arXiv.txt": {
        "abstract": "We have made an analysis of limits on the neutrino mass based upon the formation of large-scale structure in the presence of a primordial magnetic field. We find that a new upper  bound on the neutrino mass is possible based upon fits to the cosmic microwave background and matter power spectrum when the existing independent constraints on the matter density fluctuation parameter $\\sigma_8$ and the primordial magnetic field are taken into account. ",
        "introduction": "The evolution of the cosmic matter density field and large-scale structure (LSS) can be  affected by a finite  neutrino mass \\cite{Lesgourgues:2006nd}. For example, when the velocities of finite-mass neutrinos are large, the growth of density fluctuations will be impeded \\cite{1980PhRvL..45.1980B} on the scale of neutrino free-streaming. Cosmological constraints on the neutrino mass are thus of considerable interest. Recently an upper limit to the  neutrino mass of order $< 0.1-1$ eV has been deduced from a combination of tritium beta-decay endpoint experiments \\cite{2008RPPh...71h6201O} and cosmological observations \\cite{DeBernardis:2008qq,Ichiki:2008rh,Komatsu:2008hk,Ichiki:2008ye}.  At the same time, magnetic fields have been observed \\cite{Kronberg:1992pp,Wolfe:1992ab,Clarke:2000bz,Xu:2005rb} in clusters of galaxies  with a strength of $0.1-1.0~\\mu$ G. One possible explanation for such magnetic fields in galactic clusters is the existence of a  primordial magnetic field (PMF) of order  1 nG  whose field lines collapsed as the cluster formed. A variety of mechanisms for the origin and amplification of a PMF on the scale of galaxy clusters havs been proposed \\cite{Bamba:2004cu,Takahashi:2005nd,ichiki:2006sc} and the associated PMF could have influenced a variety of phenomena in the early Universe \\cite{Grasso:2000wj} such as the cosmic microwave background (CMB) \\cite{ Subramanian:1998fn, Jedamzik:1999bm, Mack:2001gc, Subramanian:2002nh, Yamazaki:2004vq, Lewis:2004ef, Kosowsky:2004zh, Yamazaki:2005yd, Challinor:2005ye, Dolgov:2005ti, Gopal:2005sg, Kahniashvili:2005xe, Kahniashvili:2006hy, Yamazaki:2006bq, Yamazaki:2006ah, Giovannini:2006kc, Yamazaki:2007oc, Kojima:2008rf, 2008PhRvD..78f3012K, Giovannini:2008aa, Kahniashvili:2008hx, Yamazaki:2009na, Kojima:2009ms, Shaw:2009nf}, and the formation of LSS \\cite{Tashiro:2005hc,Tashiro:2005ua,Yamazaki:2006mi,Sethi:2008eq,Yamazaki:2008bb,2008nuco.confE.239Y}. Of particular relevance to the present work is the fact that the effects of a finite neutrino mass on the matter density fluctuations are degenerate with the effects of a PMF. Therefore, these two effects must be constrained together when using the fluctuations of the matter density to find limits on  the  neutrino mass and the PMF parameters.  In this regard, the alternative normalization parameter $\\sigma_8$ is of particular interest as a measure of large-scale structure effects.  It is defined \\cite{Peebles:1980booka} as the root mean square of the matter density fluctuations in a comoving sphere of radius $8h^{-1}$ Mpc.  It is determined by a weighted integral of the matter power spectrum. Observations which determine $\\sigma_8$ provide information about the physical processes affecting the evolution of density-field fluctuations and the formation of structure on  cosmological scales.  The detailed mechanisms by which a PMF can affect the density field fluctuations on cosmological scales has been described in Refs.\\cite{Yamazaki:2006mi,Yamazaki:2008bb}. Of course, $\\sigma_8$ is also affected by the presence of a PMF. Indeed, in \\cite{2008nuco.confE.239Y} we demonstrated for the first time that there was a degeneracy between the parameters of the PMF and the neutrino mass.  In this article, we expand on our previous work \\cite{2008nuco.confE.239Y} by also considering the influence on the matter contributions \\cite{Kojima:2009ms,Shaw:2009nf}. By considering the effects of a PMF and finite-mass neutrinos on $\\sigma_8$ and comparing theoretically deduced values for $\\sigma_8$ with the observed range, we obtain not only insight into the underlying physical processes of density field fluctuations in the presence of a PMF, but also show that the constraint on the sum of neutrino masses from $\\sigma_8$ changes when the effects of a PMF are included. ",
        "conclusions": "Details of the effects of a PMF on the cosmological density field fluctuations can be found in Ref.~\\cite{Yamazaki:2006mi}. Here we briefly review these effects and show that a relation exists between $\\sigma_8$ and the PMF parameters. The observational constraints \\cite{Cole:2005sx,Tegmark:2006az,Rozo:2007yt,Ross:2008ze} on  $\\sigma_8$,  therefore,  lead to  a strong constraint on the PMF parameters and the neutrino mass. Since the parameters of the PMF have a strong degeneracy, the existence of a Bayesian prior constraint on $\\sigma_8$ can be used to effectively constrain the PMF. Also, since $\\sigma_8$ is constrained by diverse observational data on linear cosmological scales, we can obtain a reliable prior for use in determining the probability distributions for the parameters of the PMF from CMB observations.  The PMF effects dominate the matter power spectrum for wave numbers $k >  0.1 h$ Mpc$^{-1}$\\cite{Yamazaki:2006mi}(Fig.1).  This is because the PMF energy density fluctuations depend only on the scale factor $a$ and can survive below the Silk damping scale. Therefore, the PMF continues to be a source for the fluctuations via the Lorentz force even below the Silk damping scale. In the case of no correlation between the PMF and the matter density fluctuations, the matter power spectrum is increased by the PMF independently of whether the PMF is dominated by the pressure or the magnetic tension\\cite{Yamazaki:2006mi}.  A PMF affects the power spectrum function $P(k)$ and the matter density fluctuation $\\delta$ differently. The total density fluctuation $\\delta$ can be smaller or larger depending on whether the effect of the PMF is mainly due to  its pressure or tension. On the other hand, the power spectrum function $P(k)$ always increases in the case that the PMF does not correlate with the primordial density fluctuations.  This is  because $P(k)\\propto \\delta^2$ and is not affected by the sign of $\\delta$. \\subsection{PMF parameters and $\\sigma_8$} The alternative normalization parameter $\\sigma_8$ can be determined from a  study of the physical processes of density field fluctuations on cosmological scales within the linear regime. Recently $\\sigma_8$ has been constrained in this way by observations \\cite{Cole:2005sx,Tegmark:2006az,Rozo:2007yt,Ross:2008ze} to be in the range  $0.7 < \\sigma_8 < 0.9$. From this we can obtain strong constraints on  the PMF parameters by numerically calculating $\\sigma_8$ under the influence of a PMF.  We have shown \\cite{Yamazaki:2006bq,Yamazaki:2006mi,Yamazaki:2007oc} that the possible discrepancy between theoretical estimates and observational temperature fluctuations of the CMB for higher multipoles ($\\ell > 1000 $) could be attributed to  a PMF of comoving strength 1.0 nG$< B_\\lambda<$2.0nG. Since other constraints on the amplitude of the PMF imply $B_\\lambda \\le$10 nG\\cite{Subramanian:1998fn,Jedamzik:1999bm,Mack:2001gc,Lewis:2004ef,Kahniashvili:2008hx,Schleicher:2008dt}, our allowed range is consistent with previous works.  The value of  $\\sigma_8$ derived by a CMB analysis including such a field strength for the PMF is also $0.78 < \\sigma_8 <  0.84$. As noted above, in this work we consider $\\sigma_8$ constraints from both fits to the CMB and from large-scale structure. In this case, we adopt a prior for $\\sigma_8$ in the range of  0.75$< \\sigma_8 < $0.85 as a reasonable average of constraints from observed structure and fits to the CMB. \\subsection{Constraints on the neutrino mass and PMF parameters from $\\sigma_8$} Figure \\ref{fig2} shows the constraints on the PMF parameter $B_\\lambda$ and the sum of neutrino masses $\\Sigma m_\\nu$  for various fixed values of  $n_\\mathrm{B}$ and ranges of $\\sigma_8$ as labeled. Since the PMF power spectrum depends on $n_\\mathrm{B}$, PMF effects on the density fluctuations for small scales decrease with lower values for $n_\\mathrm{B}$. Based upon previous work \\cite{DeBernardis:2008qq,Ichiki:2008rh,Komatsu:2008hk,Ichiki:2008ye} the upper limit on the total neutrino mass is expected to be of order $< 0.1-1$ eV. For such masses the neutrino velocities can be quite  large while the mass density of neutrinos is still a significant fraction of the total dark matter density.   Therefore, the growth of density fluctuations on the free-streaming scale of neutrinos  will  be hindered.  Thus, neutrinos decrease the matter density fluctuations, while a PMF increases the matter density fluctuations (cf.~Figure \\ref{fig1}). Furthermore, for $n_\\mathrm{B}$ within the ranges constrained by previous work \\cite{Yamazaki:2004vq,Yamazaki:2006bq}, a PMF with $B_{\\lambda} ~^> _\\sim 1$ nG affects the matter density fluctuations (see Fig.~\\ref{fig2}). When the mass of the neutrinos is constrained from matter density fluctuations in the presence of a PMF, the deduced upper limit is heavier with a PMF than  without.  The expected parameters of the PMF from the combined analysis of the CMB and observed magnetic fields in galactic clusters is $B_\\lambda<$2.0nG$(1 \\sigma)$ and $<$3.0nG$(2 \\sigma)$ \\cite{Yamazaki:2004vq,Yamazaki:2006bq,Yamazaki:2009na}, while the expected value of $\\sigma_8$ based upon observations is 0.75$< \\sigma_8 < $0.85 as noted above. For this range of $\\sigma_8$, the sum of the neutrino masses is constrained to be $\\sum_{N_\\nu} m_\\nu < 0.8$ eV (from the 1$\\sigma$ constraint on  $B_\\lambda$) and $< 2.2$eV (from the 2$\\sigma$ constraint on $B_\\lambda$)  for $N_\\nu = 3$. This is a larger upper limit than that deduced previously because the effect of the PMF cancels the effect of neutrinos on the density fluctuations \\cite{DeBernardis:2008qq,Ichiki:2008rh,Komatsu:2008hk,Ichiki:2008ye}. \\subsection{Future Work} Komatsu et al. (2008) noted that the effect of a finite neutrino mass is very small for the detection range of the WMAP, e.g. $\\ell < 1000$\\cite{Komatsu:2008hk}. Therefore, as a first step toward constraining the neutrino mass together with the PMF we can use the WMAP CMB power spectrum for $\\ell < 1000$ to obtain values for the primary cosmological parameters without concern for a degeneracy between these parameters and the neutrino mass. One cannot constrain the mass of neutrinos very well from such CMB data alone. Hence, one must use other observational data which are independent of the CMB analysis to constrain the neutrino mass. For this purpose, $\\sigma_8$ is very useful since it is determined from independent observations of LSS in the regime of linear perturbation theory. Since the main purpose of the present work is to examine how the upper limit of the neutrino mass determined from $\\sigma_8$ changes when a PMF is included, we have focused on a theoretical analyses of effects of both the PMF and the neutrino mass on $\\sigma_8$ and the matter power spectrum. However, theoretical estimates of $\\sigma_8$ depend upon the cosmological model employed and $\\sigma_8$ has a significant degeneracy with $\\Omega_m = \\Omega_b + \\Omega_\\mathrm{CDM}$, $n_\\mathrm{S}$, and $A_\\mathrm{S}$, even if such primary parameters are well constrained by the CMB data, e.g. WMAP\\cite{Komatsu:2008hk}. Furthermore, there is the degeneracy between the mass of neutrinos and $\\Omega_m$ \\cite{Elgaroy:2002bi}. Therefore, one should ultimately consider the degeneracy between the PMF and other cosmological parameters including the mass of neutrinos. In subsequent work we will carry out a full-constraint on the mass of neutrinos using a Markov chain Monte Carlo (MCMC) method applied to both the CMB and the LSS data. In that effort we will need to consider the model dependence and degeneracy between cosmological parameters including the neutrino mass and PMF parameters."
    },
    "1005/1005.2688_arXiv.txt": {
        "abstract": "{Using the 3.5-m {\\em Herschel} Space Observatory, imaging photometry of Cas~A has been obtained in six bands between 70 and 500~$\\mu$m with the PACS and SPIRE instruments, with angular resolutions ranging from 6 to 37$''$. In the outer regions of the remnant the 70-$\\mu$m PACS image resembles the 24-$\\mu$m image {\\em Spitzer} image, with the emission attributed to the same warm dust component, located in the reverse shock region. At longer wavelengths, the three SPIRE bands are increasingly dominated by emission from cold interstellar dust knots and filaments, particularly across the central, western and southern parts of the remnant. Nonthermal emission from the northern part of the remnant becomes prominent at 500~$\\mu$m. We have estimated and subtracted the contributions from the nonthermal, warm dust and cold interstellar dust components. We confirm and resolve for the first time a cool ($\\sim$35~K) dust component, emitting at 70-160~$\\mu$m, that is located interior to the reverse shock region, with an estimated mass of 0.075~M$_\\odot$.} ",
        "introduction": "The large quantities of dust found in many high-redshift sources (e.g. Priddey et al. 2003, Bertoldi et al. 2003) have often been interpreted as having originated in the ejecta of core-collapse supernovae (CCSNe) from massive stars. Models for CCSNe have predicted the formation of up to 0.1-1~M$\\odot$ of dust in their ejecta (e.g. Kozasa et al 1991, Todini \\& Ferrara 2001), which could be sufficient to account for the dust observed at high redshifts (Morgan \\& Edmunds 2003, Dwek et al. 2007) and might provide a significant source of dust in the local Universe.  Cassiopeia~A (Cas~A), with an age of 330-340 years (Fesen et al. 2006) and a distance of 3.4~kpc (Reed et al. 1995), is the youngest known core-collapse SNR in the Milky Way, so the mass of swept-up interstellar material is much less than that in the ejecta. From optical spectra of distant light echoes, Krause et al. (2008) identified it as the product of a hydrogen-deficient Type~IIb CCSN. Cas~A has been intensively studied by ISO and {\\em Spitzer} at infrared wavelengths (e.g. Lagage et al. 1996; Tuffs et al. 1999; Arendt et al. 1999; Douvion et al. 2001; Hines et al. 2004; Ennis et al. 2006; Rho et al. 2008; Smith et al. 2009). Arendt et al. (1999) derived 0.038~M$_\\odot$ of 52~K dust from a fit to the {\\em IRAS} 60- and 100-$\\mu$m fluxes, while Rho et al. (2008) estimated 0.020-0.054~M$_\\odot$ of 65-265~K dust to be emitting between 5 and 70~$\\mu$m, particularly in a bright ring coincident with the reverse shock. From 450- and 850-$\\mu$m SCUBA observations, Dunne et al. (2003) reported the presence of excess emission over nonthermal flux levels extrapolated from the radio, which they attributed to 2-4~M$_\\odot$ of `cold' (T$\\sim$15-20~K) dust. However, Krause et al. (2004) argued that most of the submm excess emission could be due to dust in foreground molecular clouds and derived an upper limit of 0.2~M$_\\odot$ for cold dust within the remnant. Dunne et al. (2009) reported that the 850-$\\mu$m emission from Cas~A was polarized at a significantly higher level than its radio synchrotron emission and attributed this to $\\sim$ 1~M$_{\\odot}$ of cold dust or alternatively a significantly smaller quantity of iron needles. Iron needles were originally proposed by Dwek (2004) and produce a very different spectral energy distribution (SED) to `traditional grains', with very little flux present at $\\lambda < 500~\\mu$m. Such grains would be consistent with a high polarised fraction.  Nozawa et al. (2010) modelled the evolution of dust in Cas~A and found that the observed infrared SED of Cas~A is reproduced by 0.08~M$_\\odot$ of newly formed dust, 0.072~M$_\\odot$ of which they inferred to consist of $\\sim$40-K dust in the unshocked regions inside the reverse shock. This is supported by recent AKARI and BLAST 65-500-$\\mu$m photometric observations of Cas~A reported by Sibthorpe et al. (2010). Although they concluded that at their longest wavelengths they could not isolate any cold dust emission from the SNR from confusing interstellar emission, they did however find evidence for a $\\sim$33-K `cool' dust component, peaking at about 100~$\\mu$m, with an estimated mass of $\\sim$0.06~M$_\\odot$. In the present paper we present new far-IR and submm observations of Cas~A obtained with the {\\em Herschel} Space Observatory (Pilbratt et al. 2010). ",
        "conclusions": "Following subtraction of the nonthermal, warm dust and cold interstellar dust components, the 100-$\\mu$m image shown in Fig.~2 (top-right) shows a similar morphology to the cool dust 70-$\\mu$m image shown below it. These represent the first resolved images of this dust component, whose existence was also inferred by Tuffs et al. (2005; 60-200-$\\mu$m ISOPHOT) and Sibthorpe et al. (2010; 65-500-$\\mu$m AKARI/BLAST). The flux densities in each band from the cool dust component are listed in the final row of Table~1. They can be fitted (Fig.~3) by 0.075$\\pm$0.028~M$_\\odot$ of 35$\\pm$3-K $\\lambda^{-2}$ emissivity silicate dust having a 160-$\\mu$m absorption coefficient of 9.8~cm$^2$~g$^{-1}$ (Dorschner et al. 1995). Sibthorpe et al. derived a 33-K cool dust mass of 0.055~M$_\\odot$ (0.066~M$_\\odot$ with the dust absorption coefficients used here), consistent with our own estimate.  Nozawa et al. (2010) modelled the Hines et al. (2004) 8-100-$\\mu$m SED of Cas~A with 0.008~M$_\\odot$ of shock-heated warm dust and 0.072~M$_\\odot$ of unshocked cool dust in the remnant's interior. Their dust formation model for the Cas~A ejecta predicted 0.17~M$_\\odot$ of new dust, from which they suggested that 0.09~M$_\\odot$ had already been destroyed by the reverse shock. If the 0.075~M$_\\odot$ of cool interior dust that we find here is to survive its 5000~km~s$^{-1}$ encounter with the reverse shock, it will need to be protected by being inside very dense clumps. If most of the dust was eventually destroyed, then remnants of this type would not make a significant contribution to the dust content of the ISM, and could even dilute it.  The present observations provide no direct evidence for the presence of significant quantities of cold ($<$25~K) dust within Cas~A - the 500-$\\mu$m emission that is visible from the least obscured, northern, part of the remnant is, to first order, removed when the predicted nonthermal emission contribution is subtracted (Fig.~2; bottom-right). The cause of the 850-$\\mu$m excess in the SCUBA map of the northern part of the remnant is therefore unresolved. Spectroscopic observations are planned with PACS and SPIRE for various on- and off-remnant positions. These should enable a clean separation of line contributions and a full sampling of the continuum energy distribution from 50-670~$\\mu$m."
    },
    "1005/1005.4251_arXiv.txt": {
        "abstract": "We construct a theoretical model for low-temperature crystallization of amorphous silicate grains induced by exothermic chemical reactions. As a first step, the model is applied to the annealing experiments, in which the samples are (1)~amorphous silicate grains and (2)~amorphous silicate grains covered with an amorphous carbon layer.  We derive the activation energies of crystallization for amorphous silicate and amorphous carbon from the analysis of the experiments. Furthermore, we apply the model to the experiment of low-temperature crystallization of amorphous silicate core covered with an amorphous carbon layer containing reactive molecules.  We clarify the conditions of low-temperature crystallization due to exothermic chemical reactions. Next, we formulate the crystallization conditions so as to be applicable to astrophysical environments.  We show that the present crystallization mechanism is characterized by two quantities: the stored energy density $Q$ in a grain and the duration of the chemical reactions $\\tau$.  The crystallization conditions are given by $Q > Q\\sub{min}$ and $\\tau < \\tau\\sub{cool}$ regardless of details of the reactions and grain structure, where $\\tau\\sub{cool}$ is the cooling timescale of the grains heated by exothermic reactions, and $Q\\sub{min}$ is minimum stored energy density determined by the activation energy of crystallization.  Our results suggest that silicate crystallization occurs in wider astrophysical conditions than hitherto considered. ",
        "introduction": "Infrared observations of various comets suggest the existence of crystalline silicates in their nuclei \\citep{bregman1987, molster1999}.  In contrast, silicate in the interstellar medium is almost or entirely amorphous, namely, free from crystallites \\citep{li2007}.  It is, therefore, popular to assume that crystalline silicates observed in comets formed in the inner solar nebula  (e.g., Gail 2001; Bockelee-Morvan et al. 2002; Harker and Desch 2002; Keller and Gail 2004). This is out of harmony with the composition of the gas in cometary comae, observations of which indicate the preservation of interstellar ices in the cold outer nebula \\citep{biermann1982, mumma1996}.  As noticed by \\citet{huebner2002}, none of the thermal mechanisms proposed for silicate crystallization allows comets to retain the interstellar composition of ices in their nuclei.  Besides comets, infrared spectra tell us the presence of crystalline silicates in various kinds of object such as AGB stars, post-AGB stars, red supergiants, Herbig Ae/Be stars, and protoplanetary disks around young stellar objects. It is known that crystallization of amorphous silicate due to annealing requires temperature $T$ above 1000\\,K \\citep{hallenbeck1998, fabian2000, murata2007}.  It is generally considered that silicate does not crystallize below its glass transition temperature, which is 990 K for forsterite composition \\citep{speck2008}.  However, infrared spectra of dust shells around evolved oxygen-rich stars exhibit the presence of several emission features of crystalline silicates in their cool dust shells ($T < 300\\,$K), while silicates condense in dusty outflows with amorphous structure as evidenced by the spectra \\citep{waters1996}. Crystalline silicates are also observed in ULIRGs (ultraluminous infrared galaxies), implying that they are located in the cool, outer regions \\citep{spoon2006}.  From these observational results, several researchers claim a necessity of a yet unknown crystallization process in low-temperature environments \\citep{waters1996, molster1999, molster2001, spoon2006, watson2009}.   Physically, crystallization is a re-arrangement of atoms irregularly placed in a solid so that the atoms occupy lattice sites; in such a configuration the solid attains the lowest internal energy. However, the atoms must overcome the energy barrier, namely, activation energy of crystallization  $E$. In annealing, a fluctuation of thermal energy is used to overcome the energy barrier. However, a fluctuation of the thermal energy is not necessarily a requisite for crystallization. \\citet{carrez2002} and \\citet{kimuray08} showed that amorphous silicate with forsterite composition crystallizes at room temperature when they are irradiated by electrons in a transmission electron microscope (TEM).   Recently, we have proposed another mechanism of low-temperature crystallization due to exothermic chemical reactions of reactive molecules in an organic refractory mantles surrounding an amorphous silicate core \\citep{yamamotochigai, yamamoto2009a, yamamoto2009b}. Once the reactions are triggered by moderate heating above a few hundred kelvins, chain reactions of reactive molecules contained in the organic mantle  release heat that crystallizes the surface layer of the amorphous silicate core. We estimated the degree of crystallinity to be 0.4 to 20\\,\\% in volume for the concentration of reactive molecules of 1 to 10\\,\\%.   It was shown that the degree of crystallinity was sufficient to reproduce the observed strength of infrared features characteristic of crystalline silicates in cometary comae \\citep{kimura08, kimura2009}.  Nevertheless, we notice that the previous model of \\citet{yamamotochigai} and \\citet{yamamoto2009b} is incomplete, because the model permits crystallization at temperatures above the melting point, although, in fact, the existence of crystalline silicate is not allowed in these temperatures.  If one excludes crystallization at temperatures higher than the melting point, one would have a lower volume fraction of crystalline silicate. Furthermore, their model neglects the effect of finite particle size, namely, neglects accumulation of the heat of chemical reactions in the particle of a finite size. In consequence, we expect that the particle-size effect increases the degree of crystallinity.  These two points should be addressed in the model to properly examine the plausibility of nonthermal crystallization due to exothermic chemical reactions.   In this study, we construct a model for low-temperature crystallization of an amorphous silicate core coated  with a layer of carbonaceous material, by taking into account the melting point and  a finite particle  size  (in section 2).   In view of the analyses of cometary dust and primitive interplanetary dust particles, we consider a grain having an amorphous silicate core and a mantle of carbonaceous material.  Amorphous carbon is one of the main components in the carbonaceous material. \\citep{keller1996, keller2000, kissel1987, jessberger1999}.  We determine the activation energies of crystallization in amorphous silicate and amorphous carbon using the revised model in comparison with two previous crystallization experiments, in which particles of (1)~amorphous silicate and (2)~amorphous silicate covered with a carbonaceous layer are annealed.  These evaluations enable us to derive quantitative crystallization conditions. We demonstrate the validity of the model by its application to crystallization experiments due to exothermic chemical reactions demonstrated by \\citet{kaito2007a} in section 3.   Finally we  formulate the crystallization conditions applicable to various astrophysical environments in section 4.  We discuss the feasibility of nonthermal crystallization in astrophysical environments and summarize our conclusions in section 5. ",
        "conclusions": "We have constructed a theoretical model for low-temperature crystallization induced by exothermic chemical reactions in submicrometer-sized grains having a silicate core and a carbonaceous mantle. The validity of our model is proved by its application to the results of laboratory experiments on crystallization of amorphous silicate grains by Kaito and his colleagues. This enables us to formulate crystallization conditions of amorphous silicate in low-temperature environments using the stored-energy density $Q$ and the reaction timescale $\\tau$, irrespective of details of chemical reactions (see (\\ref{Q,tau-conditions})). It should be pointed out that the conditions given in Eq.~(\\ref{Q,tau-conditions}) are valid regardless of the grain structure, in spite of the fact that the equations are derived from the assumption of a core-mantle structure. We conclude that the low-temperature crystallization mechanism discussed in this paper should work at any dust grains unless the density $Q$ of energy stored in the grains lies below the critical density $Q_{\\rm min}$.  Our results suggest that crystallization of amorphous silicate takes place in much lower temperatures than hitherto considered. As stated in section 1, crystallization of amorphous silicate needs high temperatures ($ \\simg 1000\\,$K), while the present mechanism works by moderate heating, say, of a few hundred kelvin, depending the gas density of interest. In a protoplanetary disk, the number density of gas and its temperature at 1 AU is estimated to be $n \\sim 10^{14}\\, $cm$^{-3}$ and $T = 300\\,$K, respectively \\citep{hayashi1981}. One could derive the conditions for inducing silicate crystallization from Figs.~\\ref{fig7} and \\ref{fig8} to be $ Q_{\\rm min} = (4-8) \\times 10^{26}\\,$K\\,cm$^{-3}$ and $\\tau\\sub{cool} = 10^{-2}\\,$s. This $Q_{\\rm min}$-value corresponds to the concentration of reactive molecules in the mantle being (5-10)~\\% for the heat of chemical reaction $E\\sub{r} \\sim 10^{5}\\,$K.  This concentration of reactive molecules is in harmony with (1-10)~\\% suggested in a carbonaceous mantle of interstellar dust grains \\citep{greenberg1976}. Therefore, the nonthermal crystallization of silicate grains might come into effect in protoplanetary disks if heat of chemical reactions is deposited on the grains within a short timescale.  The minimum stored energy $Q\\sub{min}$ measures hardness of crystallization of silicate of various composition in the nonthermal crystallization; the larger $Q\\sub{min}$, the harder the crystallization is.  When the crystallization is induced by reactions (case (B) in section 4), $Q\\sub{min}$ is determined mainly by the activation energy of crystallization of the silicate (see Eq.(\\ref{condi-q}) and Eq.(\\ref{condi-q3})).  For example, the activation energies of crystallization of the amorphous silicate are $E\\sub{si}=39000$\\,K for forsterite composition and $E\\sub{si}=42000$\\,K for that of enstatite composition \\citep{fabian2000}.  The $Q\\sub{min}$ values are $ Q_{\\rm min} = 8 \\times 10^{26}\\,$K\\,cm$^{-3}$ for forsterite composition and $ Q_{\\rm min} = 9 \\times 10^{26}\\,$K\\,cm$^{-3}$ for enstatite composition.  This indicates that the crystallization of amorphous silicate of enstatite composition is harder than that of forsterite composition. On the other hand, in the case that silicate crystallization is induced by crystallization of the amorphous mantle (graphitization for an amorphous carbon mantle) due to deposition of heat of reactions in the mantle (case (A)), $Q\\sub{min}$ is determined by the activation energy of crystallization of the mantle material.  In this case,  the hardness of silicate crystallization is independent of its composition.   As long as there is a mechanism to trigger chemical reactions in a carbonaceous layer of dust grains, amorphous silicate in the grains would suffer from low-temperature crystallization. One of the plausible mechanisms is a flash heating of dust by shock waves, which has been proposed as a formation mechanism for chondrules found in meteorites. Once the dust is heated moderately in a shock region, diffusion of reactive molecules becomes active and will lead to chain chemical reactions. Note that the heating required for inducing crystallization of amorphous silicate is much lower than the one for chondrule formation. Collisions among dust particles, which raise the temperature in a local region of their surfaces ($\\sim$ tens of degree), may be another mechanism of flash heating. Since the above-mentioned mechanisms are most likely common to happen in protoplanetary disks, such a flash heating would trigger silicate crystallization in a protoplanetary disk.  Even if shock waves and mutual collisions are not effective enough to trigger chemical reactions, the low-temperature crystallization of amorphous silicate could be induced by spontaneous heating known as the Wigner storage mechanism in nuclear reactor engineering \\citep{shabalin2003}. It is well known for graphite irradiated by neutrons that energy stored in the form of lattice defects will be released spontaneously when the number density of the defects reaches a certain value. A similar phenomenon of spontaneous stored-energy release was observed in experimental studies of neutron-irradiated ices \\citep{carpenter1987, shabalin2003, kulagin2004}. \\citet{carpenter1987} found that the temperature of a cold solid methane moderator rises rapidly and uncontrollably when the heat loss becomes less effective after irradiation of fast neutrons into the moderator. The rapid temperature rise is due to the phenomenon that the energy stored at low temperatures in the form of reactive molecules was abruptly released from the irradiated methane by thermally activated diffusion of the molecules and their subsequent reactions.  Reactive molecules in the ice mantle of grains would form by irradiation of UV or high energy particles such as cosmic ray in the outer cold region of a protoplanetary disk or in a molecular cloud. When reactive molecules in such a grain are accumulated to a critical number density, the heat released by the activation of reactions would induce crystallization of amorphous silicate in the grain. However, there is no observational evidence for the presence of crystalline silicate in the outer cold region of a protoplanetary disk nor in a molecular cloud. This implies either that there is an observational selection effect to obscure the presence of crystalline silicate in those environments or that the temperatures are too low to activate chain chemical reactions in the outer cold region of a protoplanetary disk and in a molecular cloud.  If the latter is the case, observations with high spatial resolutions would constrain the conditions to trigger silicate crystallization and the accumulation rate of reactive molecules in low-temperature astrophysical environments."
    },
    "1005/1005.1642_arXiv.txt": {
        "abstract": "The existence of X-ray luminous gaseous coronae around massive disc galaxies is a long-standing prediction of galaxy formation theory in the cold dark matter cosmogony. This prediction has garnered little observational support, with non-detections commonplace and detections for only a relatively small number of galaxies which are much less luminous than expected. We investigate the coronal properties of a large sample of bright, disc-dominated galaxies extracted from the \\gimic\\ suite of cosmological hydrodynamic simulations recently presented by Crain et al.  Remarkably, the simulations reproduce the observed scalings of X-ray luminosity with $K$-band luminosity and star formation rate and, when account is taken of the density structure of the halo, with disc rotation velocity as well. Most of the star formation in the simulated galaxies (which have realistic stellar mass fractions) is fuelled by gas cooling from a quasi-hydrostatic hot corona. However, these coronae are more diffuse, and of a lower luminosity, than predicted by the analytic models of White \\& Frenk because of a substantial increase in entropy at $z \\sim 1-3$. Both the removal of low entropy gas by star formation and energy injection from supernovae contribute to this increase in entropy, but the latter is dominant for halo masses $M_{200} \\lesssim 10^{12.5}\\Msun$. Only a small fraction of the mass of the hot gas is outflowing as a wind but, because of its high density and metallicity, it contributes disproportionally to the X-ray emission. The bulk of the X-ray emission, however, comes from the diffuse quasi-hydrostatic corona which supplies the fuel for ongoing star formation in discs today. Future deep X-ray observations with high spectral resolution (e.g. with \\textit{NeXT/ASTRO-H} or \\textit{IXO}) should be able to map the velocity structure of the hot gas and test this fundamental prediction of current galaxy formation theory. ",
        "introduction": "The rarity of X-ray detections of hot gaseous coronae surrounding disc galaxies poses a fundamental challenge to the current view of how galaxies form. Whilst the possibility that galaxies might contain extra-planar gas was already raised by \\citet{Spitzer_56}, it was \\citet{White_and_Rees_78} who first proposed that hot gaseous coronae are an integral part of the galaxy formation process.  Their two-stage theory posited that hot gas reservoirs build up as gas condensing onto dark matter haloes is heated by thermodynamic shocks and adiabatic compression. Subsequent radiative cooling of this gas establishes a `cooling flow' that fuels ongoing star formation.  The ubiquity of extended soft X-ray emission from the cooling of galactic gaseous coronae was first predicted by \\citet[][hereafter WF91]{White_and_Frenk_91}, whose model built upon the ideas sketched by White \\& Rees and established the foundation for interpreting galaxy evolution within the cold dark matter cosmogony. In their framework, the dark matter haloes that host present-day disc-dominated $L^\\star$ galaxies engender gas density profiles that are conducive to efficient radiative cooling, via line-emission and thermal Bremsstrahlung, out to radii beyond the optical extent of the central galaxy. The associated cooling rate is sufficient to fuel ongoing star formation in disc galaxies, and thus offset the disruption of discs by mergers \\citep{Walker_Mihos_and_Hernquist_96,Barnes_98} and bar-instabilities \\citep{Efstathiou_Lake_and_Negroponte_82,Mo_Mao_and_White_98,Syer_Mao_and_Mo_99}, enabling the model to reproduce the abundance of morphological types observed in the local Universe \\citep[e.g.][]{Driver_et_al_06_short,Driver_et_al_07,Fukugita_et_al_07_short,Bernardi_et_al_09,Parry_et_al_09}.  The radiation associated with cooling from gas in the haloes of $L^\\star$ galaxies is predicted to fall primarily in the soft X-ray band and to have typical surface brightnessess that are readily observable with the \\xmm\\ and \\chandra\\ telescopes. This prediction is compelling because analytic and semi-analytic models based upon the WF91 framework successfully reproduce a broad range of galaxy properties, such as the luminosity function in the optical \\citep{Cole_et_al_94,Kauffmann_et_al_99,Somerville_and_Primack_99,Cole_et_al_00}, infrared \\citep{Lacey_et_al_08} and submillimetre \\citep{Baugh_et_al_05} wavebands; the bimodality in the colour-magnitude plane; the apparent `downsizing' of galaxy formation; and the black hole scaling relations \\citep{Bower_et_al_06,Croton_et_al_06,De_Lucia_et_al_06}.  The existence of X-ray luminous coronae around massive {\\em elliptical} galaxies is well established \\citep[e.g.][]{Forman_Jones_and_Tucker_85,Kim_Fabbiano_and_Trinchieri_92,O'Sullivan_Forbes_and_Ponman_01,David_et_al_06,Sun_et_al_07,Jeltema_Binder_and_Mulchaey_08,Sun_et_al_09}. However, this is generally not regarded as a test of the WF91 model because this difffuse X-ray emission could be associated with gas returned to the ISM through stellar evolution \\citep[i.e. supernovae and AGB stars, e.g.][]{Read_and_Ponman_98,Mathews_and_Brighenti_98} or with hot gas confined by the potential of galaxy groups - the most common environments of massive ellipticals \\citep[e.g.][]{Ponman_et_al_94,Trinchieri_Kim_and_Fabbiano_94,Mulchaey_et_al_96,Trinchieri_Fabbiano_and_Kim_97}. The search for diffuse soft X-ray emission from hot coronae around massive, isolated disc galaxies is therefore the most direct test of the canonical galaxy formation picture.  Adopting this premise, \\citet[][hereafter B00]{Benson_et_al_00} analysed \\rosat\\ observations of three luminous, nearby galaxies, of which two (NGC 2841, NGC 5529) are disc-dominated, whilst the third (NGC 4594, the `Sombrero galaxy') is a bulge-dominated Sa. They found no convincing evidence for diffuse X-ray emission: their upper limits on the soft X-ray luminosity are over an order of magnitude below the luminosity predicted by the WF91 cooling flow model. Similarly, an updated analysis of \\chandra\\ observations of NGC 5746 by \\citet{Rasmussen_et_al_09} failed to produce a significant detection. The greater sensitivity of \\xmm\\ and \\chandra\\ relative to \\rosat\\ has yielded the detection of a diffuse X-ray component for NGC 4594 \\citep{Li_et_al_06,Li_Wang_and_Hameed_07}, but even in this case the luminosity shortfall relative to the cooling flow model is approximately two orders of magnitude. Moreover, energy input from supernovae has been claimed as the most likely source of the emission in this case. Detections of diffuse emission from disc-dominated galaxies are, however, now becoming more common \\citep[e.g.][]{Strickland_et_al_04,Wang_05,Tullmann_et_al_06,Owen_and_Warwick_09,Sun_et_al_09} and, in general, exhibit a correlation with the star formation rate. A small number of spectacular cases exhibit X-ray emission that has an obvious biconical morphology driven by nuclear starbursts. This has led to the common interpretation that heating by supernovae is the dominant (and perhaps sole) mechanism by which local disc galaxies generate extra-planar X-ray emission, rather than the thermalisation of gravitational potential energy.  Another potential complication is the emission originating from point sources (X-ray binaries in particular), which can contribute a significant fraction of the total X-ray flux.  However, this is much less of a problem for the current generation of X-ray satellites, which have excellent angular resolution (\\textit{Chandra} in particular) that enables identification and masking of bright point sources, and high spectral resolution, which allows the removal of the contribution of unresolved, faint point sources \\citep[but see][]{Revnitvsev_et_al_08}. Typically, the removal of unresolved point sources is achieved by including a power-law component, in addition to the thermal component, when fitting the X-ray spectrum.  The dearth of convincing X-ray detections of hot coronae is puzzling and suggests that the canonical view of disc galaxy formation is incomplete. Analytic models such as WF91 necessarily rely on simplification of the complex and non-linear processes that are the essential elements of galaxy formation, such as metal-dependent gas cooling, anisotropic gas dynamics, star formation, feedback and chemical enrichment. Potentially, a more general view of the formation of hot coronae and the growth of galaxies is offered by direct, hydrodynamical simulations.  \\citet{Toft_et_al_02} highlighted a discord between hydrodynamic simulations and analytic models, finding that efficient radiative cooling can lead to the mass of \\textit{hot} gas confined by dark matter haloes being substantially lower than the predictions of cooling flow models.  Since X-ray emissivity is sensitive to gas density ($\\propto \\rho^2$), the overall X-ray luminosity of the galaxy is commensurately reduced; their model produced galactic coronae with soft X-ray luminosities two orders of magnitude lower than predicted by simple cooling flow models, and consistent with the limited observations available at the time. Interestingly, this result is in contrast to the na\\\"ive expectation of analytic models where, in general, it is assumed that cooling efficiency and X-ray luminosity are closely correlated.  It is worth noting that \\citet{Rasmussen_et_al_09} recomputed the evolution of the systems presented by \\citet{Toft_et_al_02}, using an updated simulation code and higher-resolution initial conditions. The more recent study reported X-ray luminosities that were a factor of $\\sim 2$ lower than found by \\citet{Toft_et_al_02}. Whilst these studies highlight the value of hydrodynamic simulations as a means to test the validity of the simplifications made in analytic models, their differences point to the fact that predictions derived from simulations are also subject to significant uncertainties stemming from resolution effects and the algorithmic implementation of the underlying physics.  In this paper, we investigate the X-ray halo problem using the very large \\textit{Galaxies-Intergalactic Medium Interaction Calculation} \\citep[\\gimic;][hereafter C09]{Crain_et_al_09_short}, a suite of high resolution hydrodynamic resimulations of regions drawn from the Millennium simulation \\citep{Springel_et_al_05_short}. A central aim of this project is to study the interaction of galaxies with their gas haloes and with the external intergalactic medium (IGM). The \\gimic\\ simulations are well suited to our purposes, not only because they feature many well-resolved $L^\\star$ galaxies with stellar discs, but also because their convergence behaviour is well understood (see C09), and because the only parameter that was tuned is the mass loading of winds, chosen to ensure a reasonable match to the observed cosmic star formation rate density.  The highest resolution realisations of the simulations represent the stellar discs of galaxies similar to the Milky Way with $\\sim 10^5$ particles, whilst the intermediate resolution realisations do so with a factor of 8 fewer. In both cases, the resolution is sufficient to classify the morphology of the galaxies as being either disc- or spheroid-dominated, enabling us to subsample the simulated galaxy population so as to most closely match observational samples. Combined, the intermediate-resolution simulations follow a comoving volume of $1.6\\times 10^5\\,(\\nospacehMpc)^3$, yielding approximately 460 galaxies at $z=0$ with stellar mass greater than $10^{10}\\hMsun$. The simulations include mass and energy feedback due to supernovae, and chemical enrichment from type Ia and type II supernovae and asymptotic giant branch (AGB) stars. Radiative cooling is computed on an element-by-element basis under equilibrium conditions in the presence of a realistic photoionising UV/X-ray background.  This allows us to investigate whether X-ray emission is generated primarily by the presence of a gravitationally supported coronae, or by extra-planar gas heated by supernovae. We are also able to assess the importance of metal cooling to the emission.  This work is part of the programme of the Virgo consortium for cosmological simulations. The simulations adopt the same cosmological parameters as the Millennium simulation \\citep{Springel_et_al_05_short}: $\\Omega_{\\rm m} = 0.25$, $\\Omega_\\Lambda = 0.75$, $\\Omega_{\\rm b} = 0.045$, $n_{\\rm s}=1$, $\\sigma_8 = 0.9$, $H_0 = 100~h~{\\rm km~s}^{-1}~{\\rm Mpc}^{-1}$, $h=0.73$. Our choice of cosmological parameters is consistent with the most recent determinations from the cosmic microwave background radiation and other cosmological tests, with the exception of $\\sigma_8$, which is higher by $\\sim2\\sigma$. We do not expect this difference to affect the results and conclusions of this work, since for the halo mass range considered here it results in only a small change in the epoch of formation.  The paper is laid out as follows. In \\S~\\ref{sec:methods} we describe our methods, giving a brief overview of the \\gimic\\ simulations and our main postprocessing techniques. In \\S~\\ref{sec:comp_with_obs} we confront the simulations with observational constraints, whilst in \\S~\\ref{sec:theory} we explore the analytic theory and compare it with the results of the hydrodynamic simulations. We summarise and discuss the results in \\S~\\ref{sec:summary}. A short appendix is included, in which we discuss the sensitivity of our numerical results to resolution. ",
        "conclusions": "\\label{sec:summary}  The existence of hot, X-ray luminous gas reservoirs surrounding present day Milky Way-like galaxies is a key prediction of the galaxy formation model developed by \\citet[][WF91]{White_and_Frenk_91}. This prediction is important because the WF91 model, and the many analytic and semi-analytic models based on elements of it, successfully reproduce a broad range of observed galaxy properties. However, evidence in support of this particular prediction has, in general, not been forthcoming. In this paper, we have used the \\textit{Galaxies-Intergalactic Medium Interaction Calculation} \\citep[][C09]{Crain_et_al_09_short} to investigate this challenge to the standard picture of galaxy formation using altogether different techniques. About half the galaxies that form in these five simulations are disc-dominated (an interesting fact in its own right which we will explore in detail in a further paper). Of these, 458 have stellar masses greater than $10^{10}\\hMsun$ and make up the sample we have investigated here.  \\citet[][B00]{Benson_et_al_00} searched for the predicted extended emission around three promising candidate disc galaxies using the \\rosat\\ telescope, but they failed to detect any.  Subsequent extensive surveys with the \\xmm\\ and \\chandra\\ telescopes finally succeeded in detecting diffuse emission in the soft (0.5-2.0\\keV) X-ray band around some disc galaxies, but at much lower levels than predicted by WF91. The larger number of non-detections in these surveys has tightly constrained upper limits on this kind of emission. The observational data now offer a clearer picture than was available to B00 who were the first to highlight the potential implications of the absence of X-ray emission for galaxy formation theory. In sum, the data indicate that the soft X-ray luminosities of disc galaxies are one to two orders of magnitude lower than predicted and that they broadly correlate with both the $K$-band luminosity and the star formation rate of the galaxy.  The \\gimic\\ hydrodynamic simulations make different assumptions and approximations to those of the WF91 model. However, we find that these simulations also generate X-ray coronae around disc galaxies, although these are much weaker that predicted by WF91 for the reasons summarized below. Many of the bulk X-ray properties of the simulated galaxies agree with observations. In particular, the 458 galaxies in our sample exhibit similar scaling relations to the data. For this comparison, we used a compilation of X-ray measurements and upper limits from several studies \\citep{Benson_et_al_00,Strickland_et_al_04,Wang_05,Tullmann_et_al_06,Sun_et_al_07,Owen_and_Warwick_09,Rasmussen_et_al_09} and extracted $K$-band luminosites from the \\twomass\\ database and disc rotation velocities (derived from inclination-corrected 21-cm measurements) from the \\textsc{HyperLeda} database. We find that the galaxies in the simulations broadly reproduce the observed scaling and scatter of the $L_{\\rm X}-L_{\\rm K}$, $L_{\\rm X}-v_{\\rm rot}$ and $L_{\\rm X}-\\dot{M}_\\star$ relations at $z=0$.  There are two main reasons why the X-ray emission from the \\gimic\\ galaxies is much weaker that predicted by WF91. Firstly, the mass of hot gas within the galaxies' haloes is, typically, lower than in the WF91 model. This is partly because some of the gas has been consumed into stars and partly because some gas has been blown out in winds. Secondly, and most importantly, the distribution of the hot gas is much less concentrated than the distribution of dark matter, violating one of the main assumptions of the WF91 model. The gas is more extended than expected because its entropy has been significantly raised. For galaxies like the Milky Way and fainter, the increase in entropy is mostly due to the injection of energy generated by supernovae. The gas is placed in a high adiabat at early times, $z\\sim 1-3$, when the star formation peaks.  The entropy of the corona can also be affected by radiative cooling. Cooling acts to reduce the gas temperature by selectively removing the lowest entropy gas and locking it into stars, thus raising the mean entropy of the remaining coronal gas \\citep[e.g.][]{Voit_et_al_02}. The effect of radiative cooling in this respect is therefore not dissimilar to that of outflows. However, we have shown that the reduction of the hot gas density by radiative cooling is only significant for the most massive galaxies in our sample. In this regime, however, our simulations overestimate the effects of radiative cooling because they do not include additional energetic feedback processes such as AGN. This shortcoming is manifest, for example, in the high stellar fractions of the largest galaxies which are about a factor of 2 larger than inferred observationally \\citep{Guo_et_al_09}. By contrast, the median stellar fractions of the Milky Way-like galaxies and smaller in our simulations are in very good agreement with these data, although the scatter in the simulations appears perhaps to be too large. Simulations that include the effects of AGN are required to understand the X-ray properties of the gas associated with the most massive galaxies.  The main difference between the hydrodynamic simulations and the semi-analytic models based on the WF91 formalism is that the simulations follow the kinematics and dynamics of the gas self-consistently. Although outflows are often included in semi-analytic models \\citep[see e.g.][]{Kauffmann_99,Cole_et_al_00,Bower_et_al_01,Baugh_et_al_05, Bower_et_al_06,Croton_et_al_06_short,Bertone_DeLucia_and_Thomas_07}, the readjustment of the gas to energy injection is not generally taken into account. One exception is the recent work by \\citet{Bower_McCarthy_and_Benson_08} which attempts to follow self-consistently the heating of the intracluster medium by AGN and its effects on the subsequent cooling of the gas. The key contribution of the full hydrodynamic treatment is the ability to follow the way in which the entire structure of the gas corona adjusts in order to accommodate the increase in entropy associated with the energy injection by supernovae (or AGN). It is important to note, however, that the \\gimic\\ simulations assume a specific model for energy feedback. Although this appears well motivated, it is essential to investigate how the X-ray properties of the gas depend on the implementation of energy feedback processes. Whilst methods such as that developed by \\citet{Bower_McCarthy_and_Benson_08} provide a blueprint for starting to explore the reaction of gas to energy injection in a computationally inexpensive fashion, we note that model-independent conclusions can only be arrived at with suites of hydrodynamical simulations featuring a wide range of possible feedback treatments (an approach adopted by the \\owls\\ simulations, albeit at lower resolution than required here). We intend to pursue this methodology with high-resolution resimulations of individual galaxies in future studies.  In spite of the approximations inherent in the analytic treatment of gas dynamics in the WF91 model, key results of this model are borne out by the \\gimic\\ simulations. The predicted quasi-hydrostatic hot coronae around $L^\\star$ disc galaxies do indeed form, as anticipated, by the shock heating of infalling gas as the dark matter halo develops. In all but a few low mass galaxies, emission from these coronae dominates the galactic diffuse X-ray luminosity.  Galactic winds contribute to the X-ray luminosity but, in general, their emission is subdominant. Winds represent a small fraction of the coronal gas, but since they tend to be hotter, denser and more metal-rich than the quasi-hydrostatic component, they contribute disproportionally to the X-ray luminosity. Since most X-ray observations target the central parts of disc galaxies, where the wind gas is concentrated, and since the observed X-ray luminosity correlates well with the star formation rate, it is easy to misinterpret the importance of winds to the overall X-ray emission. In fact, in our simulations the X-ray luminosity also correlates with the star formation rate, just as in the data, even though the emission is dominated by the quasi-hydrostatic corona. This correlation occurs because both the star formation rate and the total amount of hot gas depend on the mass of the halo. The simulations do produce a few examples of bright galaxies in which outflowing gas contributes the bulk of the X-ray luminosity. These may be counterparts of starburst galaxies such as M82 \\citep[whose intense star formation seems to be associated with interaction with M81, see e.g.][]{Chynoweth_et_al_08}. In our simulations, star formation in $L^\\star$ disc galaxies today is mostly fuelled, as found by WF91, by the cooling of gas out of the hot corona. This can occur either through a smooth cooling flow \\citep[e.g.][]{Thomas_et_al_86} or through fragmentation into high-velocity clouds \\citep[e.g.][]{Maller_and_Bullock_04,Booth_and_Theuns_07}.  Our simulations indicate that hot outflows do not, in general, dominate the overall soft X-ray luminosity of isolated disc galaxies. We therefore expect gravitational inflow to dominate over supernova-driven outflows, and the majority of a galaxy's soft X-ray luminosity to stem from hot, quasi-hydrostatic gas at low surface brightnesses. This physical picture is, in principle, testable with future X-ray observatories with instrumentation of high sensitivity and spectral resolution, such as \\textit{NeXT/ASTRO-H} and \\textit{IXO}, because coronal gas dynamics follow directly from X-ray line diagnostics. We therefore expect these facilities directly to verify, or rule out, the key r\\^ole in the formation of disc galaxies that theoretical models assign to hot gaseous coronae."
    },
    "1005/1005.3523_arXiv.txt": {
        "abstract": "To ascertain whether magnetic dynamos operate in rocky exoplanets more massive or hotter than the Earth, we developed a parametric model of a differentiated rocky planet and its thermal evolution.  Our model reproduces the established properties of Earth's interior and magnetic field at the present time.  When applied to Venus, assuming that planet lacks plate tectonics and has a dehydrated mantle with an elevated viscosity, the model shows that the dynamo shuts down or never operated.  Our model predicts that at a fixed planet mass, dynamo history is sensitive to core size, but not to the initial inventory of long-lived, heat-producing radionuclides.  It predicts that rocky planets larger than 2.5 Earth masses will not develop inner cores because the temperature-pressure slope of the iron solidus becomes flatter than that of the core adiabat.  Instead, iron ``snow'' will condense near or at the top of these cores, and the net transfer of latent heat upwards will suppress convection and a dynamo.  More massive planets can have anemic dynamos due to core cooling, but only if they have mobile lids (plate tectonics).  The lifetime of these dynamos is shorter with increasing planet mass but longer with higher surface temperature.  Massive Venus-like planets with stagnant lids and more viscous mantles will lack dynamos altogether.  We identify two alternative sources of magnetic fields on rocky planets: eddy currents induced in the hot or molten upper layers of planets on very short period orbits, and dynamos in the ionic conducting layers of ``ocean'' planets with $\\sim$10\\% mass in an upper mantle of water (ice). ",
        "introduction": "Earth distinguishes itself in the inner Solar System with a magnetic dipole generated by a dynamo in its iron core.  Mars had a magnetic field of comparable intensity that collapsed early in the planet's history \\citep{Arkani04} and never re-appeared \\citep{Lillis08}. Mercury has a weak (1\\% of Earth) global field \\citep{Stanley05,Christensen06b}.  Venus currently lacks a dynamo, and we have no evidence for or against one in its past \\citep{Nimmo02}. At least one massive satellite in the outer Solar System (Ganymede) appears to have a core dynamo \\citep{Schubert96}.\\\\  Planets as small as twice the mass of the Earth have been discovered on short-period orbits around other stars \\citep{Mayor09}.  A fraction of these transit the parent star, allowing mean densities to be determined.  The smallest transiting exoplanet has a mass of $4.8\\pm0.8$ \\mearth~and a mean density of $5.6\\pm 1.3$ g~cm$^{-3}$, consistent with an Earth-like composition \\citep{Leger09,Queloz09}. The second smallest ($6.6\\pm0.9$ \\mearth) has a mean density only one-third that of Earth, indicating it has a thick volatile envelope \\citep{Charbonneau09}. These two planets may rotate synchronously and have effective emitting temperatures of around 2600~K and 500~K, respectively (that of Earth is 255~K).  The {\\it Kepler} spacecraft mission \\citep{Koch10,Borucki10} is expected to find many such ``hot'' massive rocky planets \\citep{Selsis07,Gaidos07}.\\\\  Do these planets have magnetodynamos?  While there have been many theoretical studies of the dynamos in the Earth and smaller Solar System bodies, as well as in the ice and gas giants \\citep{Breuer09,Stevenson09}, rocky planets more massive than the Earth have not been considered, although simple dynamo scaling laws have been applied to gas giants on short-period orbits \\citep{Griessmeier04}.  Magnetic fields may protect planetary atmospheres against erosion by stellar winds and coronal mass ejections \\citep{Griessmeier04,Lundkin07,Dehant07}.  Interaction between a giant planet's magnetic field and that of the host star has been inferred from corotational chromospheric activity \\citep{Shkolnik08,Walker08} and may produce detectable radio emission \\citep{Lecavelier09}.  Ohmic dissipation of the kinetic energy of winds in partially ionized atmospheres has been proposed as an explanation for the ``inflated'' radii of some short-period gas giants \\citep{Batygin10}.\\\\  For a planet to have a dynamo, it must contain a fluid layer that is sufficiently electrically conducting for the magnetic Reynolds number Re$_m$ = $V L/\\lambda > 40$, where $V$ and $L$ are a characteristic fluid velocity and length, respectively, $\\lambda = 1/(\\mu_0 \\sigma$) is the magnetic diffusivity, $\\mu_0 = 4\\pi \\times 10^{-7}$~N~A$^{-2}$, and $\\sigma$ is the electrical conductivity.  In Earth-like planets the conducting fluid is liquid iron or an alloy thereof ($\\lambda \\sim 2$ m$^2$ sec$^{-1}$ \\citet{Stevenson09}).  The possibility that the oxidation state of planetary material might prevent iron core formation was explored by \\citet{Elkins08}, but we assume that such cores do form.  Metallization of silicates and miscibility with Fe is not expected until a pressure of 20~TPa \\citep{Oganov05}.  In the cores of Earth-size planets ($L \\sim 3 \\times 10^6$ m), convective motions as small as $10^{-4}$~m~sec$^{-1}$ are sufficient for Re$_m$ $> 40$.  The temporal variation of Earth's magnetic field, such as westward drift, imply velocities of $10^{-4}-10^{-3}$~m~sec$^{-1}$ at the top of the core.\\\\  A dynamo also requires a source of energy to maintain convective motions against internal ohmic dissipation.  Core convection can be driven by the removal of heat from the core faster than it can be transported by conduction, by the release of latent heat during phase transitions (i.e., solidification of iron), or by the formation of buoyant fluid via the expulsion of light elements during iron solidification.  The last two are related to the first because the latent heat of solidification must be removed.  Thus dynamo activity is ultimately related to the transport of heat and the temperature contrast across the core-mantle boundary (CMB).  That contrast can be maintained either by efficient cooling of the lower mantle (e.g., by descending slabs of former lithosphere during plate tectonics) or by the heating of the core by radioactive nuclides.  The absence of a dynamo in Venus has been ascribed to the lack of plate tectonics, which slows cooling of the mantle, as well as the core \\citep{Nimmo02}.\\\\  In dynamo theory, the condition for convection is expressed as a requirement on entropy production, rather than energy production, because the latter is unaffected by dynamo dissipation when mechanical energy is converted to heat.  The balance between entropy sinks and sources in the core gives the maximum dissipation $\\Phi$ by the dynamo.  If $\\Phi > 0$ then a dynamo is permitted.  We calculate $\\Phi$ using parameterized models of the interior structure and thermal evolution of a planet consisting of an Fe core and a silicate mantle.  We account for, but do not vary, the presence of light elements in the core.  We estimate the magnetic Reynolds number and average surface field using previously established scaling laws.  We address whether dynamos operate in rocky planets with different masses, surface temperatures, and modes of mantle convection.  Figure \\ref{fig.schematic} is a schematic of our simplified planet.  In the equations that follow, we employ the following subscripts: 1 (ambient), p (planet or surface), m (mantle), c (core or core-mantle boundary), i (inner core boundary), and 0 (center).\\\\ ",
        "conclusions": "}  \\subsection{Interpretation of trends and implications}  For an inner core to form, enough sensible heat must be removed for the center temperature to reach the solidus.  We obtain an approximate timescale $t_c$ for this by setting $Q_s = -Q_K$ and performing the appropriate integration.  This gives \\begin{equation} \\label{eqn.growthtime} t_c \\approx \\frac{M_c\\Delta S}{8\\pi kR_c} \\left(1 - \\frac{D^2}{R_c^2}\\ln \\frac{\\tau_{0}({\\rm Fe})}{\\tau_{c}({\\rm pv})}\\right). \\end{equation} where the solidii for Fe and perovskite are at central and CMB conditions, respectively.  For $k=40$~W~m$^{-1}$~K$^{-1}$, and $R_c \\sim M_p^{1/4}$ (Figure \\ref{fig.massradius}), the first factor is $\\approx 2 M_p^{3/4}$~Gyr.  The second factor is of order unity: Both solidii increase with planet mass and so we expect $t_c$ to increase slightly faster with planet mass than $M_p^{3/4}$.  This is born out by our detailed calculations, i.e a 2 $M_{\\oplus}$ planet takes more than twice as long to form an inner core than an Earth-mass planet (Figure \\ref{fig.masstempmode}).  A maximum time $t_s$ for the complete solidification of the entire core, after the inner core starts growing, is obtained by dividing the total latent heat by the heat flow.  (In our cases, the heat flow through the mantle tends to be close to or not much larger than the conduction through the core itself.)  In this case, $t_s$ is dependent on mass only through the pressure-dependences of certain material properties of the core: \\begin{equation} \\label{eqn.freezetime} t_s \\approx \\frac{c_p  \\Delta S}{2 k \\alpha G}. \\end{equation} For standard values, $t_s \\sim 50$~Gyr. The iron cores of rocky planets with Earth-like composition will never completely solidify, and solid inner cores never include more than about half of the total core mass in 10~Gyr.  Thus, planets that do form inner cores ($\\le$2 \\mearth) will have dynamos that last longer than the main-sequence life of the parent star.  In planets with masses of 2-3 \\mearth, inner cores do not grow but instead iron condenses as ``snow'' at the CMB, shutting down convection and the dynamo within 10~Gyr.  In yet more massive planets, the core never cools sufficiently for iron to condense.  These planets have persistent dynamos produced by core cooling.  However, the temperature contrast between the core and the mantle decays, and the heat flow across the CMB eventually falls to the core conduction value.  At that point, $\\phi = 0$ and the dynamo ceases.\\\\  At a given planet mass, the surface field $B_p$ increases with increasing core size, and a solid inner core formation forms earlier (Figure \\ref{fig.sizeandrad}a).  The former is a consequence of the cubic distance attenuation law and a shallower CMB, as well as increased heat flow across a thinner mantle.  The latter is a result of more rapid cooling, but also higher pressure and a higher solidus at the planet's center.  Dynamo history is much less sensitive to the initial abundance of long-lived radionuclides in the planet's mantle (Figure \\ref{fig.sizeandrad}b) because of the self-regulating nature of viscosity-dependent mantle convection \\citep{Tozer72}.  $B_p$ generally decreases with planet mass.  Equations \\ref{eqn.bfield2} and \\ref{eqn.bsurf} show why: The intensity of the field in the core $B_c$ depends on the available entropy production per unit mass $\\phi$ and only weakly on core size.  All else being equal, $\\phi$ will decrease with mass because the core of a larger planet has less surface area per unit mass and a thicker mantle through which to lose heat.  If $\\phi \\sim R_c^{-1}$, this makes $B_c$ insensitive to planet mass.  The intensity of the field at the surface will decrease as $R_c^{-3}$ and thus will roughly scale as $M_p^{-3/4}$.\\\\  Higher surface temperature leads to elevated mantle temperatures and lower viscosities.  Although the former lowers the temperature contrast across the CMB, the latter more than compensates because of the exponential decrease of viscosity with temperature.  The outcome, increased magnetic field strength and dynamo lifetime, is clearly seen in Figure \\ref{fig.masstempmode}.\\\\  \\subsection{CoRoT-7b}  Our model predicts that CoRoT-7b has a dynamo at its estimated age of 1.3-2.2~Gyr \\citep{Leger09}, but only if it has a mobile lid.  In this case, the dynamo is driven entirely by core cooling, and no inner core is present.  In fact, it is the absence of iron solidification (as ``snow'') that allows a dynamo to persist.  Whether or not plate tectonics operates on CoRoT-7b is debated \\citep{Valencia07,ONeill07,Kite09}, and the presence of a dynamo may be one means of testing this.  However, a core dynamo is not the only mechanism that can produce a magnetic field on such an object (see \\S \\ref{sec.other}).\\\\  \\subsection{Uncertainties and neglected effects \\label{sec.uncertainties}}  {\\it High-pressure material properties:} The properties of silicates and Fe at pressures of hundreds or thousands of GPa are very uncertain.  Although equations of state become dominated by electron degeneracy pressure at pressures well above 1~TPa, other parameters, namely $k$ and $\\alpha$, are very important in determining whether a dynamo operates.  \\citet{Aitta06} has proposed that iron has a tricritical point at around 800 GPa at which separate solid and liquid phases disappear and the latent heat of fusion vanishes.  These conditions are reached at the center of a 2.5~\\mearth~planet and at the CMB of a 6~\\mearth~planet (Figure \\ref{fig.massradius}).\\\\  {\\it Mantle structure:} We assume a fully convecting mantle: A layered mantle will support a higher temperature contrast for the same heat flow \\citep{McNamara00}, and this would tend to suppress dynamo operation.  The lowest 150~km of the Earth's mantle (within the D'' layer) may be in the stability field of a post-perovskite (ppv) phase \\citep{Oganov04} and ppv may dominate in sufficiently cool ``super-Earth'' mantles \\citep{Valencia09}.  If a magma ocean existed at the base of the Hadean terrestrial mantle, release of latent heat during crystallization would have suppressed the cooling of Earth's core \\citep{Labrosse07b}.  Our calculations indicate that the mantles of planets with M $\\ge$ 2.5\\mearth~and stagnant lids are partially molten.  We do not attempt to model the thermal evolution of such objects but they are (at least initially) unlikely to support a dynamo.\\\\  {\\it Other effects of light element in the core:} Our model does not include melting-point depression in the core by light elements, particularly sulfur (S).  The effect of these is to delay inner core formation relative to the pure Fe case.  In that sense our calculations are optimistic in predicting dynamos.  (This is in contrast to the case of small bodies where S can maintain a dynamo by delaying the freezing of the entire core).  However, we choose our initial core temperature based on the present size of the Earth's inner core and thus, relative to Earth, this calibration may cancel some of the effect of S.  We include the production of buoyancy flux resulting from the exclusion of light elements from a growing solid core.  However we assume that light elements are uniformly mixed into the outer liquid core and do not include possible variations in light element concentrations with planet mass.  The liquid silicate-liquid metal partition coefficient of sulfur increases with pressure \\citep{Li03} but no experimental data is yet available for pressures exceeding those at Earth's CMB.  If the trend continues to higher mass (and for other elements), the power available for dynamos in the cores of super-Earths may be smaller than our model predicts.\\\\  {\\it Non-radiogenic sources of mantle heat:} Dissipation of tidal strain could be an additional source of heat in the mantles of non-synchronously rotating planets close to their parent stars \\citep{Barnes10}.  The timescale for rotational synchronization is a few Myr \\citep{Guillot96} while that of orbital circularization can be Gyr \\citep{Rasio96}, but both are strongly dependent on semimajor axis.  Sufficiently intense heating during this period could lead to a long-lived thermal anomaly in a mantle that influences the transport of heat much like a stagnant lid.  We also ignore the potential complication of heating by late giant impacts \\citep{Robert09,Arkani10}.\\\\  {\\it Effect of the stellar magnetic field:} In the absence of an ionosphere, the ambient magnetic field of the host star $B_*$ could reach the planet's core: A sufficiently strong applied field will inhibit convective motion in transverse directions.  This will occur when the Stuart number \\begin{equation} \\label{eqn.stuart} {\\rm St} = \\frac{\\sigma B_*^2 R}{\\rho v} \\gg 1. \\end{equation} Using the definition of Re$_m$ we re-write this as: \\begin{equation} \\frac{\\sigma B_*^2 R^2}{\\rho \\lambda} \\gg {\\rm Re_m}. \\end{equation} For typical values $\\sigma \\sim 5 \\times 10^5$~S~m$^{-1}$, $\\rho \\sim 10^4$~kg~m, $\\lambda \\sim 2$~m$^2$~sec$^{-1}$, $L \\sim 3 \\times 10^6$~m, the condition becomes $B_* \\gg 60$Re$_m^{1/2}$ nT.  For Re$_m \\sim 10^3$ the requirement becomes $B_* \\gg 2$~$\\mu$T.  $B_* \\sim$ 100~$\\mu$T near the surface of the Sun: Assuming that the stellar magnetic field decreases with the square of the distance (because of interplanetary plasma), $St \\gg 1$ only within 0.1~AU.  Thus inhibition of convection could occur in the cores of the very closest planets, although the increased temperature and conductivity of surface rocks at such proximity to the star may block the stellar field (see \\S \\ref{sec.other}).\\\\  \\subsection{Other sources of magnetic fields \\label{sec.other}}  In a planet is embedded in an ambient magnetic field from its star that varies with orbital period $P$, eddy currents will generate an induced field if a layer is sufficiently thick ($d$) and conducting. The condition for a substantial induced field is \\citep{Stevenson03} \\begin{equation} d\\sigma \\gg \\frac{1}{2\\pi \\mu_0 R_p}. \\end{equation} For an Earth-size planet where $d$ is in km and $P$ in days, the required conductivity $\\sigma$ is $\\gg 2 P d^{-1} (R_p/R_{\\oplus})^{-1}$~S~m$^{-1}$.  The conductivity of anhydrous molten silicates is 0.1-1~S~m$^{-1}$ while that of molten carbonates is three orders of magnitude higher \\citep{Gaillard08}.  A planet with a magma ocean only a few km deep on a planet with a very short-period orbit such as CoRoT-7b (0.85~d) could have an induced field.\\\\  Uranus and Neptune have magnetic fields that are thought to be generated in a layer of high pressure water (ice) \\citep{Cavazzoni99}. The ionic conductivity of planetary ices is about $2 \\times 10^3$~S~m$^{-1}$ at pressures above 40~Gpa, conditions that are reached at 0.7-0.8 radii in these planets \\citep{Stanley06}.  We added liquid water and high-pressure ice (Ice VII) layers to the interior model described in \\S \\ref{sec.interior} to calculate the pressure at the base of an upper water/ice mantle as a function of total planet mass and water mass fraction (Figure \\ref{fig.ice}).  A basal pressure of 40~Gpa is reached on planets of a few $M_{\\oplus}$ and ice fractions of $\\sim$ 10\\%.  Such an envelope of water could be easily accomodated by the observed radius (2.7 $R_{\\oplus}$) of GJ~1214b \\citep{Charbonneau09} but not by the mass and radius of CoRoT-7b \\citep{Valencia10}.\\\\  \\subsection{Conclusions}  Our model of a planet consisting of a silicate mantle and pure Fe core reproduces the gross interior structure and thermal history of the Earth.  Using the magnetic field strength scaling of \\citet{Aubert09}, the predicted evolution of the field strength is also consistent with paleomagnetic inferences and the present value.  When applied to Venus, with a stagnant lid, an elevated surface temperature, and enhanced mantle viscosity, our model corrctly predicts the absence of a dynamo at present.  We are thus confident in extrapolating the model to rocky planets with approximately Earth-like composition but with different mass, surface temperature, and/or mode of convection.\\\\  The model predicts that 1-2 \\mearth~planets with plate tectonics will sustain a dynamo for up to 7~Gyr by core cooling before an inner core forms and the dynamo intensifies at a time that depends on surface temperature, mass, and the relative size of the core.  Planets in this mass range with stagnant lids will have dynamos only once an inner core begins to grow.  Planets with higher surface temperatures can more easily sustain a dynamo because the resulting higher mantle temperatures lead to a lower viscosity and hence more rapid mantle convection and core cooling.\\\\  Planets larger than about 2 \\mearth~do not develop solid inner cores because the adiabat is steeper than the solidus.  In such planets, iron freezes first at or near the core-mantle boundary.  The latent heat and any light elements released at the top of the core will not be able to drive a long-lived dynamo.  In planets between 2-3 \\mearth, this leads to the premature shutdown of any dynamo within 10~Gyr.\\\\  In planets with masses $>$ 3 \\mearth, iron never begins to freeze in 10~Gyr.  If these planets have plate tectonics, our model predicts that they have relatively weak dynamos driven solely by core cooling but which can last several Gyr, depending on mass and surface temperature.  The mantles of planets with stagnant lids in this mass range are hot enough to prevent core cooling and dynamo operation, but they may also be partially molten.  This last scenario is not addressed by our model.  Our model predicts that CoRoT-7b may have a dynamo, but only if the planet has a mobile lid, and the strength of the field at its surface will be less than that at of the present Earth ($\\sim 50 \\mu$T).  Advances in our knowledge of the properties of planetary materials at very high pressure will improve our ability to make more quantitative predictions about the intensity and longevity of core dynamos in massive rocky exoplanets.\\\\"
    },
    "1005/1005.1526_arXiv.txt": {
        "abstract": "{ M\\,82 is a unique representative of a whole class of galaxies, starbursts with superwinds, in the Very Nearby Galaxy Survey with Herschel. In addition, its interaction with the M\\,81 group has stripped a significant portion of its interstellar medium from its disk. SPIRE maps now afford better characterization of the far-infrared emission from cool dust outside the disk, and sketch a far more complete picture of its mass distribution and energetics than previously possible. They show emission coincident in projection with the starburst wind and in a large halo, much more extended than the PAH band emission seen with Spitzer. Some complex substructures coincide with the brightest PAH filaments, and others with tidal streams seen in atomic hydrogen. We subtract the far-infrared emission of the starburst and underlying disk from the maps, and derive spatially-resolved far-infrared colors for the wind and halo. We interpret the results in terms of dust mass, dust temperature, and global physical conditions. In particular, we examine variations in the dust physical properties as a function of distance from the center and the wind polar axis, and conclude that more than two thirds of the extraplanar dust has been removed by tidal interaction, and not entrained by the starburst wind.} ",
        "introduction": "\\object{M\\,82} is the closest specimen of starbursts with superwinds, at a distance of only $(3.9 \\pm 0.3)$\\,Mpc \\citep{Sakai99}. Its almost edge-on orientation is ideal for the study of the superwind, and more generally the interaction of galaxies with the intergalactic medium and the role of the gas and dust dynamics in the self-regulation of starbursts.  The exceptional star formation activity of M\\,82 is thought to be caused by its tidal interaction with the M\\,81 galaxy group, redistributing the gas within the disk of M\\,82 toward the central regions where it accumulates, and expelling atomic gas in giant streams connecting the disk to the intragalactic medium \\citep{Yun94}. M\\,82 is an evolved starburst, in the sense that its activity has been sustained (episodically) over much of its lifetime. Coupled stellar population synthesis and photoionization models reveal two major episodes of star formation over the past 10\\,Myr \\citep{Forster03}. Numerous young supernova remnants have been detected \\citep{Muxlow94}. Another peak of star formation about 150\\,Myr ago is inferred from the age-dating of stellar clusters \\citep{Konstantopoulos09}, and 200 super star clusters, signposts of starbursts, have been found in M\\,82 \\citep{Melo05}, spanning a wide age interval from 60\\,Myr to several Gyr \\citep{Gallagher99, Grijs01}.  \\cite{Ward03} found that more than half the total molecular gas of the central regions is warmer than 50\\,K, and \\cite{Mao00} found that CO transitions up to $J=7-6$ trace diffuse gas abundantly illuminated by UV radiation. \\cite{Panuzzo10} modeled CO transitions from $J=4-3$ to $J=13-12$ in the central 43$^{\\prime\\prime}$, observed with the SPIRE FTS spectrometer, and they suggest that turbulent heating caused by stellar winds and supernovae may be the dominant excitation process for the highest transitions, tracing gas at about 500\\,K. The molecular gas properties thus reinforce the notion that M\\,82 is an evolved starburst, which has been active for long enough to completely disrupt the cold component of the interstellar medium.  The mechanical energy simultaneously released by a large collection of massive stars and supernovae is able to drive a superwind detected in both X-rays emitted by shocked gas \\citep{Strickland97} and optical lines excited by a combination of photoionization and shock ionization \\citep{Shopbell98}. Optical emission around the wind also partly issues from a reflection nebula \\citep{Scarrott91}, indicating the presence of dust. Abundant molecular gas is associated with both the superwind (although probably not spatially coincident but rather outside the bicylindrical outflow) and tidal streams \\citep{Walter02}, and PAH emission is found in a very extended halo \\citep{Engelbracht06}.  This paper studies the spatially-resolved properties of the cool dust detected by SPIRE. Since PACS photometric observations had not been scheduled at the time of writing, and Spitzer MIPS images are saturated, we sample the far-infrared spectral energy distribution only longward of 250\\,$\\mu$m, where it can be safely assumed that the emission is dominated by grains in thermal equilibrium. SCUBA maps have higher angular resolutions than SPIRE maps; their sensitivities are however too low to reliably detect any emission outside the disk of M\\,82, even when all archival data are combined \\citep{Leeuw09}. In contrast, bright filaments and diffuse emission are detected out to very large distances with SPIRE. We thus focus on the dust content and physical properties outside the main body of M\\,82, where SPIRE maps provide truly unique information. ",
        "conclusions": "The intergalactic medium can be enriched in metals by both tidal stripping and starburst winds. M\\,82 offers an exquisite example of both processes at work in the same galaxy. Considering the spatial distributions of the dust brightness and temperature in the halo, the tidal interaction may be far more effective than the wind in removing the cool interstellar medium from the disk of M\\,82. We estimate that more than two thirds of the extraplanar dust mass is associated with the tidal streams traced by atomic hydrogen. If dust is able to survive for a long time in the extragalactic medium around such interacting systems, which may have been common in the past, then it will have significant column densities along lines of sight to distant galaxies, and cause extra attenuation outside the Milky Way."
    },
    "1005/1005.4629_arXiv.txt": {
        "abstract": "{We point out an interesting theoretical prediction for elliptical galaxies residing inside galaxy clusters in the framework of modified Newtonian dynamics (MOND), that could be used to test this paradigm. Apart from the central brightest cluster galaxy, other galaxies close enough to the centre experience a strong gravitational influence from the other galaxies of the cluster. This influence manifests itself only as tides in standard Newtonian gravity, meaning that the systematic acceleration of the centre of mass of the galaxy has no consequence.  However, in the context of MOND, a consequence of the breaking of the strong equivalence principle is that the systematic acceleration changes the own self-gravity of the galaxy. We show here that, in this framework, initially axisymmetric elliptical galaxies become lopsided along the external field's direction, and that the centroid of the galaxy, defined by the outer density contours, is shifted by a few hundreds parsecs with respect to the densest point.} ",
        "introduction": "Modified Newtonian Dynamics \\cite{Milgrom1983a,BM1984} is a recipe linking the gravitational field in galaxies with their baryonic content. This recipe makes many successful predictions on galaxy dynamics, such as the baryonic Tully-Fisher relation \\cite{McGaugh2005}, and the precise shape of rotation curves of both high and low surface-brightness galaxies in the field \\cite{McGaugh_deBlok1998}, as well as those of tidal dwarf galaxies (e.g., Gentile et al. 2007). It could also hold clues to the observed universality of dark and baryonic gravities at the effective dark halo core radius \\cite{Donato_etal2009,Gentile_etal2009,Milgrom2009a}. The successes of this recipe in galaxies could be an emergent phenomenon, linked with the complex feedback between baryons and dark matter in the process of galaxy formation \\cite{Peebles2009}, but a more radical explanation of these successes is a modification of gravity on galaxy scales. This interpretation requires to postulate additional dark matter at extragalactic scales (e.g., Sanders 2007; Angus et al. 2007; Angus 2009), but let us note that most relativistic versions of such modified gravity theories inevitably include dark fields that seed structure growth, and may play the role of dark matter in galaxy clusters \\cite{Ferreira_Starkman2009}. This reduces calculability on large scales, and means that the modified gravity interpretation of the MOND recipe is more easily falsifiable on galaxy scales. Falsifying such an alternative approach to the problems of galactic dynamics is actually needed in order to more firmly anchor the current cold dark matter (CDM) paradigm as the one that should ultimately be confronted with data at the sub-galactic, galactic and cosmological scales.  One of the unique properties of the modified gravity version of MOND in galaxies, compared to its Newtonian CDM counterpart, is that it violates the Strong Equivalence Principle when considering the external gravitational field in which a system is embedded. Hence, the effects of the gravitational environment are important, and fundamental concepts such as the escape speed from a system are fully dependent on the external field. This can lead to extremely interesting and testable predictions. For instance, \\cite{Milgrom2009b} analytically predicted in the Solar System an anomaly due to the external field of the Galaxy, in the form of a quadrupole field that may be detected in future measurements of planetary and spacecraft motions. On the other hand, \\cite{FBZ2007} and \\cite{Wu_etal2007,LMC} studied the effect of the external field from large scale structure on the Milky Way and showed that the MOND predictions for the galactic escape speed were quite realistic, while allowing the Large Magellanic Cloud to be on a bound orbit.  Galaxies inside clusters present a very interesting environment where the external field has a significant influence. These galaxies experience both their self-gravity and the gravity of the other galaxies of the cluster (except for the galaxy at the centre of the cluster): the latter manifests itself only as tides in standard Newtonian gravity, meaning that the systematic acceleration of the centre of mass of the galaxy has no consequence.  But in the context of MOND, a consequence of the breaking of the strong equivalence principle is that the systematic acceleration changes the own self-gravity of the galaxy.  In this Letter, we emphasize that, in MOND, the combination of the self-gravity with the external acceleration breaks the front-back symmetry of a purely self-gravitational system. As elliptical galaxies are abundant in clusters, they provide an ideal case to test this front-back asymmetry.  Offsets such as those discussed in this letter have also been studied in the recent literature in the context of MOND N-body simulations \\cite{phantom_peak}, CDM N-body simulations \\cite{LZK2009}, and lensing observations \\cite{SFFZ2008}. ",
        "conclusions": "We estimate that for typical parameters of elliptical galaxies residing in large clusters, a centroid shift of a few hundreds of parsecs is expected in MOND, as compared to the centroid of the outer isodensity contours of the galaxy.  While tides can also cause some lopsidedness in Newtonian and MONDian gravity, there is a major distinction with the pure MOND effect under scrutiny here. The tidal effects are small inside the tidal radius, and are traditionally neglected when dealing with dense elliptical centres. The fact that these effects can be neglected in our analysis is confirmed by the fact that the average density of our model galaxy inside the saddle point where the external field starts to dominate is larger than the cluster density. The saddle point is estimated to be 15 kpc from the Plummer sphere (total mass $10^{11}\\Msun$) centre (Fig.1  of \\cite{phantom_peak}).  Here the average density enclosed is $\\sim 4\\times 10^{-3}{\\rm M}_\\odot {\\rm pc}^{-3}$, typically more than one order of magnitudes larger the typical cluster density at that same radius ($1000 \\times \\rho_{\\rm crit} \\sim 10^{-4} {\\rm M}_\\odot {\\rm pc}^{-3}$). Note that \\cite{ZT2006} argues that such tidal criteria holds in MOND as well in Newton.  Thus, we conclude that:  \\begin{itemize}  \\item If real elliptical galaxies in rich clusters show perfectly symmetric light with {\\it no} significant offsets between the centroids of the inner and outer contours, the classical version of MOND is likely excluded. This predicted lopsidedness of galaxies inside distant rich clusters should be falsifiable with photometry {\\it only}, using, e.g. the VLTI. To resolve a centroid offset of ~200pc in an elliptical galaxy in a rich cluster of typical internal gravity $\\sim 10^{-8} \\,$cms$^{-2}$  at a distance of 160-210 Mpc (e.g., Abell 1983, Abell 2717, MKW9, \\cite{Pointecouteau_etal2005,Sanders2003} would require a minimum angular resolution of 0.5 arcsec.  The offsets at the centroids of galaxies are between 10 pc and 500 pc: for an elliptical in Coma (100 Mpc far away) this corresponds to 0.02 to 1 arcsec, while for an elliptical in Virgo (20 Mpc far away) it corresponds to 0.1 to 5 arcsec.  \\item On the other hand, observing such a lopsidedness would {\\it not} consitute a direct falsification of the CDM paradigm. Indeed, CDM simulations predict that isolated CDM halos are typically lopsided due to a lack of relaxation, so this should be the case inside clusters too. \\cite{Maccio_etal2007} studied the shapes of a large sample of ten thousand pure CDM halos, and found that any offset between the barycentre (CoM) and the density-centre (potential minimum) is correlated with a massive satellite inside an unrelaxed halo, i.e., the offset is a measure of unrelaxedness in the CDM context. This is characterized by an offest of order of 1 or 2 percent between the central dark matter density maximum and the centre of mass at the virial radius. Let us note that if this can be translated into a predicted lopsidedness for the light distribution of isolated ellipticals in CDM, this would be a clean test since MOND does not predict such an offset (except at extremely large radii) for isolated galaxies. What is more the centres of some ellipticals are home to strong star formation which could also distort the observed central isophotes. The isophotal distortions predicted in MOND, if actually observed, could thus be due to this. However, for galaxies inside clusters, these more mundane offsets would not be expected to be aligned with the direction of the local external field pointing to the cluster centre. If a statistically relevant preferred direction is found for the lopsidedness of galaxies residing in clusters, it would thus be in support of MOND, although it would not constitute a direct falsification of CDM. \\end{itemize}  Nevertheless, a caveat is that our simple models here were based on a constant external field of $\\sim 10^{-8} \\,$cm~s$^{-2}$. However, in reality, while galaxies are orbiting inside clusters, the external field acting on them varies with time, both in direction and amplitude. So the next step will be to make the external field vary as a function of time in our simulations, to check how this affects the predicted lopsidedness.  Finally, we also note that the flattened and lopsided potential created by the external field would also be valid in disk galaxies, although most of them are in the field or at the outskirts of galaxy clusters. Still, the flattened MOND potential would create a differential force with a component normal to the disk, hence a specific torque. This could cause differential precession of the disk angular momentum vector, and could lead to the formation of warps, even in a very low external field for isolated galaxies~\\cite{CT2009}.  {\\bf Acknowledgements:} The work has been performed under the Project HPC-EUROPA (211437), with the support of the European Community - Research Infrastructure Action under the FP8  ``Structuring the European Research Area\" Programme. We thank Luca Ciotti, Pasquale Londrillo, Carlo Nipoti for generously sharing their code. XW acknowledges hospitality of the Sterrewacht of Leiden University and of the AIfA of Bonn University. XW acknowledges additional support from HPC-EUROPA (211437). XW is supported by a scholarship of the Scottish Universities Physics Alliance (SUPA).  BF is a CNRS Senior Research Associate. HSZ acknowledges partial support from Dutch NWO fellowship and Royal Society visitorship."
    },
    "1005/1005.4842_arXiv.txt": {
        "abstract": "Motivated by the fact that cosmological perturbations of inflationary quantum origin were born Gaussian, the search for non-Gaussianities in the cosmic microwave background (CMB) anisotropies is considered as the privileged probe of non-linear physics in the early universe. Cosmic strings are active sources of gravitational perturbations and incessantly produce non-Gaussian distortions in the CMB. Even if, on the currently observed angular scales, they can only contribute a small fraction of the CMB angular power spectrum, cosmic strings could actually be the main source of its non-Gaussianities. In this article, after having reviewed the basic cosmological properties of a string network, we present the signatures Nambu--Goto cosmic strings would induce in various observables ranging from the one-point function of the temperature anisotropies to the bispectrum and trispectrum. It is shown that string imprints are significantly different than those expected from the primordial type of non-Gaussianity and could therefore be easily distinguished. ",
        "introduction": "\\label{sec:introduction}  The origin of cosmic strings dates back to the discovery that cosmological phase transitions triggered by the spontaneous breakdown of the fundamental interaction symmetries may form topological defects~\\cite{Kirzhnits:1972,Kobsarev:1974, Kibble:1976}. Cosmic strings belong to the class of line-like topological defects, as opposed to point-like monopoles and the membrane shaped domain walls. As shown by Kibble, the appearance of defects in any field theory is related to the topology of the vacuum manifold~\\cite{Kibble:1976}. If the ground state of a field theory experiences a spontaneous breakdown from a symmetry group $G$ to a subgroup $H$, Kibble showed that cosmic strings will be formed if the first homotopy group $\\pi_1(G/H) \\neq I$ is non-trivial. In other words, if non-contractile loops can be found in the manifold $\\calM=G/H$ of equivalent vacua. Similarly, the other homotopy groups $\\pi_0$ and $\\pi_2$ determine the formation of domain walls and monopoles respectively. Once formed and cooled, these defects cannot be unfolded, precisely due to their non-trivial topological configuration over the vacuum manifold of the theory. This simple statement suggests that cosmic strings, and topological defects in general, are a natural outcome of the unification of the fundamental interactions in the context of Cosmology. As remnants of unified forces, their discovery would be an incredible opportunity to probe extremely high energy physics with ``a telescope''.  In the last thirty years, many works have been devoted to the cosmological consequences, signatures and searches for topological defects~\\cite{Hindmarsh:1994re, Vilenkin:2000, Sakellariadou:2006qs, Peter:1208401}. They have pushed cosmic strings to the privileged place to be generically compatible with observations. Indeed, domain walls and monopoles are prone to suffer from the cosmological catastrophe problem: their formation is sufficiently efficient (or their annihilation sufficiently inefficient) to either overclose the universe or spoil the Big-Bang Nucleosynthesis (BBN) predictions~\\cite{1996RPPh.59.1493S, 2005APh.23.313C}. For domain walls, this implies that either they should be extremely light, i.e. formed at an energy scale less than a few $\\MeV$, or no discrete symmetry should have been broken during the cooling of the universe. There is not so much choice for the monopoles: if interactions were unified, monopoles would have been formed. The homotopy group of $\\pi_2(G/H)$ with $H$ containing the $U(1)$ of electroweak interactions is indeed non-trivial\\footnote{As often with topological defects, sensitivity to the underlying model is such that one can often find a counter-example of any result. Both of these statements, on walls and monopoles, can be evaded in some particular models or with some amount of fine-tuning, as for instance if cosmic strings can be attached to them and catalyse annihilations~\\cite{PhysRevD.26.435, PhysRevLett.45.1}.}. Cosmic inflation was originally designed to solve the monopole problem. If a phase of accelerated expansion of the universe occurs, then any defects will be diluted enough to no longer have any (dramatic) consequences on cosmology~\\cite{Guth:1980zm, Starobinsky:1980te, Linde:1981mu, Starobinsky:1982ee}. Meanwhile, Inflationary Cosmology solves the flatness and homogeneity problem of the standard Big-Bang model, explains the origin and spectrum of the Cosmic Microwave Background (CMB) anisotropies, as the formation of the large scale structures~\\cite{Mukhanov:1990me, Martin:2006rs}. Inflation provides a priori an easy solution to the topological defects problem by diluting them to at most one per Hubble radius. However, one has to keep in mind that this mechanism works only if the defects were formed before inflation, and even in that case some may survive~\\cite{2007PhRvD.76h3510A}. This has to be the case for monopoles and heavy walls, but not for local strings. On the contrary, exhaustive analysis of particle physics motivated inflationary models, embedding the Standard Model $SU(3)\\times SU(2) \\times U(1)$, has shown that strings are generically produced at the end of inflation~\\cite{Jeannerot:2003qv}. In this picture, our universe should contain cosmic strings whose properties are closely related to those of the inflaton~\\cite{Rocher:2004my, Rocher:2006nh, Battye:2006pk}. String Theory provides an alternative framework to Field Theories: brane inflationary models propose that the accelerated expansion of the universe is induced by the motion of branes in warped and compact extra-dimensions~\\cite{Becker:2007, Dvali:1998pa, Alexander:2001ks, Kachru:2003sx}. Inflation ends when two branes collide and such a mechanism again triggers the formation of one-dimensional cosmological extended objects, dubbed cosmic superstrings~\\cite{Burgess:2001fx, Sarangi:2002yt, Dvali:2003zj, Jones:2003da}. These objects may be cosmologically stretched fundamental strings or one-dimensional D-brane~\\cite{Davis:2005, Copeland:2009ga}. Although cosmic superstrings are of a different nature than their topological analogue, they produce the same gravitational effects and share similar cosmological signatures~\\cite{Sakellariadou:2008ie, Sakellariadou:2009ev}.  Among the expected signatures, cosmic strings induce temperature anisotropies in the CMB with an amplitude typically given by $GU$, where $U$ is the string energy per unit length\\footnote{To avoid any confusion with Greek tensor indices, we will use the Carter's notations $U$ and $T$ for the string energy density and tension~\\cite{Carter:1989dp}.} and $G$ the Newton constant~\\cite{Bouchet:1988}. For the Grand Unified Theory (GUT) energy scale, one has $GU \\simeq 10^{-5}$, which precisely corresponds to the observed amplitude of the CMB temperature fluctuations~\\cite{Durrer:2001cg}. However, the power spectra do not match: topological defects are active sources of gravitational perturbations, i.e. they produce perturbations all along the universe history, and cannot produce the characteristic coherent patterns of the acoustic peaks~\\cite{Durrer:1996, Magueijo:1996px, Albrecht:1997, Spergel:2006hy, Komatsu:2008hk}. Current CMB data analyses including a string contribution suggest that they can only contribute to at most $10\\%$ of the overall anisotropies on the observed angular scales~\\cite{Bouchet:2000hd, Fraisse:2007}. For Abelian cosmic strings (see Sec.~\\ref{sec:kinds}), numerical simulations in Friedmann--Lema\\^{\\i}tre--Robertson--Walker (FLRW) spacetimes show that this corresponds to an upper two-sigma bound $GU < 7\\times 10^{-7}$~\\cite{Hindmarsh:2007}. Direct detection searches provide less stringent limits but are applicable to all cosmic string models: $GU< 4\\times 10^{-6}$~\\cite{Gilbert:1995de, Jeong:2007, Amsel:2007ki}. Detecting cosmic strings in the CMB certainly requires one to go further than the power spectrum~\\cite{Gangui:2001fr, Pogosian:2007gi} (see however Sec.~\\ref{sec:cls}). In fact, strings induce line-like discontinuities in the CMB temperature through the so-called Gott--Kaiser--Stebbins effect, which are intrinsically of non-Gaussian nature~\\cite{Gott:1984ef,Kaiser:1984iv}. In the inflationary picture, cosmological perturbations find their origin in the quantum fluctuations of the field--metric system, and therefore were born generically Gaussian. Non-Gaussianities can nevertheless appear from non-linear effects during inflation or from couplings to other fields (see the other articles in this issue). These non-Gaussianities are of the primordial type, i.e. they exist before the cosmological perturbations reenter the Hubble radius. On the other hand, cosmic strings are a source of non-Gaussianity at all times and, as we will see, produce different signals from the CMB point of view\\footnote{Notice that second order perturbations, being non-linear, actively generate non-Gaussianities but at a relatively small amplitude~\\cite{Sefusatti:2007ih, Pitrou:2008ak, Nitta:2009jp, Boubekeur:2009uk}.}.  In this article, we review the non-Gaussian features a cosmological network of cosmic strings produce in the CMB anisotropies. In a first section, we briefly scan various cosmic string models and emphasize their similarities and differences for cosmology. Making observable predictions for cosmic strings faces the problem of understanding their cosmological evolution. Not only one has to solve the local dynamics in curved space, but as extended objects, cosmic strings follow a globally non-local evolution: the fate of one string depends on its interactions with the others. The cosmological evolution of a network of cosmic strings is a non-trivial problem which can be overcome by means of numerical simulations. These simulations permit an estimation of the various statistical properties affecting the observational signatures, such as the number of strings per Hubble radius, their shapes, velocities or the loop density distribution. Latest results in this area, for the Nambu--Goto (NG) type of cosmic strings, are presented in Sec.~\\ref{sec:evol}. Once the statistical properties of a cosmological cosmic strings network are known, it is possible to extract meaningful observables depending only on the unique model parameter $U$\\footnote{If no currents are flowing along the string, Lorentz invariance implies that the string tension $T$ equals the energy density $U$.}. In Sec.~\\ref{sec:smapng}, we recap the expected CMB temperature anisotropies induced by cosmic strings, derived from various methods. Particular attention is paid to small angle CMB maps which preserve all of the projected statistical information. We then derive the cosmic string signals expected in various non-Gaussian estimators ranging from the one-point function of the CMB temperature fluctuations to the bispectrum and trispectrum. We conclude in Sec.~\\ref{sec:conc} and discuss various non-Gaussian aspects which still have to be explored. ",
        "conclusions": "\\label{sec:conc}  The results presented in this article were essentially concerned with Nambu--Goto type of cosmic strings, which is the simplest realisation of a one-dimensional spatially extended object. As a result, they should not be blindly extrapolated to other types of string, although, as argued in Sec.~\\ref{sec:kinds}, some of them are expected to be generic. In particular, due to the scaling of the long strings, cosmic string loops do not influence significantly the CMB observables. Changing the intercommuting probability is expected to rescale some of the presented results~\\cite{Sakellariadou:2004wq, Avgoustidis:2007aa}, but in a way which remains to be quantified.  In Sec.~\\ref{sec:evol}, we have briefly reviewed the current understanding of the cosmological evolution of a string network by means of FLRW numerical simulations, which is a non-trivial problem even for NG strings. Observable predictions crucially depend on this step. Numerical simulations can be avoided by making some assumptions on the string distribution but at the expense of introducing unnecessary extra parameters. When approximate analytical models are then used to derive observable predictions, one should keep in mind that the results are as uncertain as the values assumed for the additional parameters. Provided one is interested in length scales not affected by gravitational back-reaction effects, all of the statistical properties of a NG string network in scaling depends only on one unknown physical parameter: the string energy density per unit length $U$, not more\\footnote{The expansion rate is supposed to be known.}.  In this context, Sec.~\\ref{sec:smapng} discusses the non-Gaussian effects induced by a cosmological string network in the CMB temperature anisotropies. We have shown that string induced CMB fluctuations have a negative skewness and a non-vanishing kurtosis. On a CMB temperature anisotropies map, these non-Gaussianities imprint characteristic signatures in a multifractal analysis as well as in the gradient magnitude, both being more significant at small angles. This property is recovered in Fourier space: the CMB angular power spectrum decays at most as $1/\\ell$, for the large multipoles $\\ell$, and strings become the dominant sources of primary fluctuations. The skewness appears to be the direct consequence of the breaking of the time reversal symmetry in an expanding universe, and implies the existence of a non-vanishing bispectrum. Using analytical approximations, tested and confirmed by numerical simulations, we then derived the expected bispectrum and trispectrum of string induced CMB temperature anisotropies for the large multipoles. Although the bispectrum decays not faster than $\\ell^{-6}$, the trispectrum multipole dependency is in $\\ell^{-\\rho}$, where $\\rho=6+1/(\\chi+1)$ and $\\chi$ is a small number related to the tangent vector correlator and the NG loop distribution. Due to the line-like CMB patterns induced by the strings, both the bispectrum and trispectrum are enhanced on all elongated triangle and quadrilateral configurations of the wavevectors. These ones may constitute the best configurations to look for a non-Gaussian string signal while being experimentally limited by finite beam resolution. Let us note that our expressions have been derived in the flat sky approximation. String non-Gaussianities at small multipoles is still an open problem which could be dealt with full sky string maps~\\cite{Landriau:2004, Regan:2009hv, Landriau:2010cb}. However, if, as the current constraints suggest, cosmic strings marginally contribute to the large scale CMB anisotropies, then they should show up at large multipoles in all of the above mentioned observables. This is precisely where the experimental efforts are directed. At very small angular scales, the difficulties will certainly be to separate the string signals from the astrophysical sources. Interestingly, the very soon accessible intermediate angles probed by the Planck satellite, and the other ground based telescopes, may not suffer from this problem and could be an open window on cosmic strings.  \\ack It is a pleasure to thank Patrick Peter, Mairi Sakellariadou, Daniele Steer and Teruaki Suyama for a careful reading of the manuscript and their enlightening comments. This work is supported by the Belgian Federal Office for Scientific, Technical and Cultural Affairs, under the Inter-University Attraction Pole grant P6/11."
    },
    "1005/1005.5243_arXiv.txt": {
        "abstract": "{ We investigate the structure of the wind in the neutron star X-ray binary system Vela~X-1 by analyzing its flaring behavior. Vela~X-1 shows constant flaring, with some flares reaching fluxes of more than 3.0\\,Crab between 20--60\\,keV for several 100\\,seconds, while the average flux is around 250\\,mCrab. We analyzed all archival \\textit{INTEGRAL} data, calculating the brightness distribution in the 20--60\\,keV band, which, as we show, closely follows a log-normal distribution. Orbital resolved analysis shows that the structure is strongly variable, explainable by shocks and a fluctuating accretion wake. Analysis of \\xte ASM data suggests a strong orbital change of \\nh. Accreted clump masses derived from the \\inte data are on the order of $5\\times10^{19}$--$10^{21}$\\,g. We show that the lightcurve can be described with a model of multiplicative random numbers. In the course of the simulation we calculate the power spectral density of the system in the 20--100\\,keV energy band and show that it follows a red-noise power law. We suggest that a mixture of a clumpy wind, shocks, and turbulence can explain the measured mass distribution. As the recently discovered class of supergiant fast X-ray transients (SFXT) seems to show the same parameters for the wind, the link between persistent HMXB like Vela~X-1 and SFXT is further strengthened. } ",
        "introduction": "The high-mass X-ray binary (HMXB) Vela~X-1 was discovered in 1967 \\citep{chodil67a} and is today one of the best studied objects of its class hosting a neutron star. It is a bright, eclipsing source, showing strong pulsations with a pulse period of around $283$\\,sec \\citep{mcclintock76a} and an orbital period of 8.9\\,days \\citep{vankerkwijk95a}. \\citet{quaintrell03a} have shown that the separation between the neutron star and its optical companion, the B0.5Ib supergiant HD\\,77581, is only $1.7$\\,stellar radii. The neutron star is thus deeply embedded in the stellar wind of the supergiant, which has a mass loss rate of about $10^{-6}\\,\\msun\\,$yr$^{-1}$ \\citep{nagase86a}. Material from the wind is accreted by the neutron star and is channeled by its strong magnetic field onto the magnetic poles. The accretion rate leads to an average X-ray luminosity of $\\sim$ $4 \\times 10^{36}$\\,erg\\,s$^{-1}$, although the luminosity is strongly variable on all time-scales, varying up to a factor of at least 20--30 \\citep{staubert04a, kreykenbohm08a}. \\citet{staubert04a} found short, but very bright flares reaching up to almost 7\\,Crab (1000\\,counts\\,sec$^{-1}$ [cps] in the 20--40\\,keV band of \\inte ISGRI).  The physical processes leading to this strong variation are unclear, however. Because the luminosity of the neutron star is proportional to the mass accretion rate, changes in the density or the effective velocity of the stellar wind directly lead to variations in the X-ray luminosity \\citep{bondi44a, davidson73a}. Changes of that kind have shown up in simulations of the system \\citep{blondin90a, blondin91a, mauche08a}. Strong density changes have also been proposed to explain varying line strengths in high-resolution X-ray spectra \\citep{sako99a, watanabe06a} or were inferred via measurement of emission lines in hot, thin plasma at the same time with fluorescence lines in colder, optical thick plasma \\citep{schulz02a, goldstein04a}.  The goal of this paper is to study the variations in the lightcurve and investigate a possible connection between the wind structure and the flaring behavior. This paper is structured as follows: in Sect. \\ref{sec:obsdat} we present our data and analysis method. In Sect. \\ref{sec:ana_all} we show an orbital phase averaged and orbital phase resolved analysis of the \\inte ISGRI hard X-ray and \\xte ASM soft X-ray lightcurve. In order to constrain the physical processes leading to the observed variability we present a numerical method to simulate lightcurves in Sect. \\ref{sec:simu}. For this simulation we calculate and model the power spectral density of Vela~X-1. In Sect. \\ref{sec:disc} we summarize our results and discuss them in a larger context. ",
        "conclusions": "\\label{sec:disc} \\subsection{Summary} We presented the first analysis of the statistical distribution of the brightness of Vela~X-1, with special regard to the flaring behavior. Our main results are: \\begin{itemize} \\item The orbital phase averaged brightness distribution above 20\\,keV can be characterized by a log-normal distribution. \\item The orbital phase resolved histograms indicate systematic variations besides the eclipse. \\item The brightness distribution in the soft X-rays is strongly affected by absorption. \\item Analysis of the soft X-rays confirm a decline in the average countrate with orbital phase. \\item The brightness distribution of a simulated lightcurve consisting of a multiplication of individual random processes  can be used to describe the data. \\end{itemize}  \\subsection{Giant flares} From our results obtained by analyzing 3.6\\,Msec of \\inte ISGRI data we can calculate the probability to measure extremely bright flares in that time range. Extremely bright flares are defined by their countrate being larger than 300\\,counts\\,sec$^{-1}$ in the 20--60\\,keV band. The calculated probability is $\\sim$0.011\\% for the best-fit Gaussian, while it is $\\sim$0.023\\% for the simulated lightcurve. These numbers agree well with $\\sim$0.02\\%  measured directly from our dataset. We conclude that bright flares are rare, but not singular events. No separate process is required for their explanation.  \\subsection{Wind structure and accretion geometry} The origin of the constant flaring behavior lies in the accretion flow onto the neutron star, which is obviously not smooth, but highly structured. Neglecting any absorption and scattering effects, the luminosity distribution gives us the opportunity of calculating the mass distribution in this flow in units of mass per time. To calculate the mass accretion rate it is necessary to know the absolute luminosity of the X-ray source in the given energy band. As the luminosity is strongly energy-dependent, we modeled the energy spectrum with a power law with Fermi-Dirac-cutoff, using the parameters of \\citet{kreykenbohm08a}, who found these parameters for data from revolutions 137--141, which were part of our analysis. The overall spectrum of our data could be described with these values, too. With those parameters and a distance of 2\\,kpc \\citep{nagase86a} the median absolute luminosity $\\left<L_\\text{x}\\right>$ is $ 5.1\\times10^{36}$\\,erg\\,sec$^{-1} $.  During the accretion process only part of the potential energy is converted to X-rays. We assume an accretion efficiency of $\\eta = 0.3$ and calculate the median accretion rate \\begin{equation}\\left<\\dot M\\right> =  \\frac  {\\left<L_\\text{x}\\right>}  {c^2 \\cdot \\eta} = 3.0\\times10^{-10}\\,\\text{M}_\\odot\\,\\text{yr}^{-1} \\label{eq:mdot} \\end{equation} The multiplicative standard deviation is $\\tilde \\sigma_{\\dot M}  = 1.9$ in the 20--60\\,keV energy range. Figure \\ref{fig:histo_mdost.eps} shows the distribution of inferred $\\dot M$ values and the corresponding luminosity. To show the log-normal behavior of the distribution more clearly, we used a linear scaled $x$-axis. We included data not only from the 20--60\\,keV energy band, but also from the softer 20--30\\,keV band and the harder 40--60\\,keV band in this figure. All three bands show the same behavior. Slight deviations between the different energy bands are due to the rough spectral fit, because only the limited energy range of ISGRI was available for fitting. Nonetheless it is evident that all three curves follow a log-normal distribution in a range of $\\dot{M}$ very well expected for this kind of object regarding the mass loss rate of the optical companion and the assumed accretion efficiency.  \\begin{figure} \\centering \\includegraphics[width=0.95\\columnwidth]{histo_alltimes_10_mdot.pdf} \\caption[Histogram of mass accretion rate]{Histogram of the calculated accretion rate $\\dot{M}$ of the neutron star. The rate was calculated using a standard spectrum and the countrate in the 20--60\\,keV band (black), 20--30\\,keV band (red), and 40--60\\,keV band (green). For detailed description see text. } \\label{fig:histo_mdost.eps} \\end{figure}  \\subsubsection{A clumpy wind} If we assume that the neutron stars accretes directly from the wind, we can infer the mass distribution in the stellar wind from the accretion rate. Strong density variations in the wind comparable to the variations in the accretion rate can be described by a model of a clumped stellar wind, a model based on instabilities in the the line-driven acceleration mechanism \\citep[see, e.g.,][]{feldmeier03a, dessart05a, oskinova07a}. A clumpy wind is also supported by observations. Based on \\textsl{ASCA} data, \\citet{sako99a} show that the stellar wind of HD\\,77581 is most likely strongly clumped by measuring line strengths in the spectra and showing that these lines are only in accordance with a mass loss rate of $\\sim$$10^{-6}\\,M_\\odot\\,$yr$^{-1}$ when assuming a clumped wind, consisting of dense, cold clouds and a hot, ionized surrounding medium. The clouds make up for about 90\\% of the wind mass, while the hot ionized medium makes up for about 95\\% of the volume, i.e., the volume filling factor of the clumps is $f_v=0.05$. \\citep{sako03a}. Further evidence for clumping in the stellar wind was brought forth by different authors using high-resolution \\chandra spectra \\citep{schulz02a, goldstein04a, watanabe06a}. In these spectra emission lines from H- and He-like ions indicate a hot thin gas. The spectra, however, showed additionally strong fluorescent lines from lower ionized atoms, e.g, of Fe, Mg, and Ne. These originate in colder, optically thick plasma, which can be explained as clumps in the hot surrounding gas. Additionally, if a dense clump is passing through the line of sight between the observer and the X-ray emitting region, the absorption coefficient will rise significantly.  Clumps measured through absorption dips \\citep{nagase86a} may very well be the same clumps that are accreted and are thus directly responsible for luminosity fluctuations measurable up to at least 100\\,keV \\citep{sako99a}.  The accretion of a clump or a similar dense feature in the wind will lead to an increased X-ray luminosity of the source, i.e., a flare. As the lightcurve of Vela~X-1 is dominated by those flares, which are strongly overlapping, it is hard to estimate the length of any individual flare. Visual inspection of the lightcurve indicates, however, that the shortest flares last around $t_\\mathrm{fl} \\approx 2.5$\\,ksec. During that time $\\langle M \\rangle_\\text{acc} = \\langle\\dot{M}\\rangle t_\\mathrm{fl} = 4.7\\times10^{19}$\\,g of material will be accreted onto the neutron star, assuming the average accretion ratio from Eq.~\\ref{eq:mdot}. Using the $\\beta$-law \\citep{castor75a}, the average density at the neutron star orbit gives values of $\\rho_\\mathrm{wind} \\approx 8\\times10^{-18}\\,\\mathrm{g}\\,\\mathrm{cm}^{-3}$. In their simulations \\citet{blondin90a,blondin91a} showed that regions with an overdensity of $\\sim$100 are possible in the filaments of the accretion wake. Assuming a spherical clump with such an increased density, i.e., a volume filling factor $f_v=0.01$ \\citep{fullerton06a}, a radius of $r_\\mathrm{cl} = 2.37\\times10^{10}$\\,cm is obtained. These clumps would show up in absorption with a maximal column density of $N_\\mathrm{H}^\\mathrm{cl} \\approx 9\\times 10^{21}\\,\\mathrm{cm}^{-2}$, i.e., close to the observed absorption column \\citep{kreykenbohm08a}, especially considering that more than one clump could be in the line of sight at any given moment and that we assumed short flares, and therefore small clumps.  Besides the  normal flares, giant flares are also visible in the lightcurve and are more easily constrained in time. The largest flare in the investigated lightcurve is at MJD 52971.2, lasting for 38\\,ksec with an average count rate of $\\sim$140\\,cps. During that time  $M_\\text{gf} \\approx 2\\times 10^{21}$\\,g were accreted, using the same conversion factor and efficiency as in Eq.~\\ref{eq:mdot}. This mass is in the regime of the values given in numerous other works, see for example \\citet{negueruela08a} or \\citet{walter07a}. Note, however, that these authors have investigated supergiant fast X-ray transients (SFXT), a rather new class of X-ray sources. In a model proposed by \\citet{negueruela06a}, these SFXT are regarded as HMXB, similar to systems like Vela~X-1, with the orbit parameters being the major difference. This similarity extends to similar companion stars, and thus similar wind properties. The similar wind properties are also evident in the irregular flaring behavior of the SFXT IGR J08408$-$4503, which was explained by a clumpy wind \\citep{leyder07a}. This explanation was also given for the flaring behavior of ``regular'' HMXB \\citep[see, e.g.,][]{vandermeer05a} and is supported by our analysis. We conclude that the calculated clumpsizes and masses for SFXT can be very well compared with the values of large clumps of Vela~X-1.  \\subsubsection{Temporal smearing} In the previous section we showed that the clump distribution is probably roughly log-normal. If true, this result contradicts to typical assumptions for mass distributions in the atmospheres of high-mass stars, where power-law distributions are generally assumed \\citep[e.g.,][]{kostenko99a}. But it is unlikely that the distribution of clumps in the stellar atmosphere translates directly to a distribution of accretion rate from the wind, because temporal smearing is taking place during accretion. If the clump mass distribution in the wind would be indeed log-normal, smearing could lead to small deviations from this log-normal distribution in the measured brightness distribution. We tried to analyze the effect of smearing by folding the simulated lightcurve with an exponential function with a characteristic timescale. Due to the log-normal character of the lightcurve, strong smearing leads to a higher average countrate, which disagrees with the measured lightcurve distribution. We therefore included another fit parameter, allowing the renormalization of the simulated lightcurve. With this method we could find no reliable timescale for the smearing, because our time resolution was limited to 283.5\\,sec and all emerging fit values were of the order of or lower than that resolution. The analysis of a lightcurve with higher time resolution of 20\\,sec, but with a moving average of 283.5\\,sec applied to smooth out the intrinsic pulse variations came to the same result: the smearing timescale is around the level of the time resolution and thus not significant. From this simulation we can conclude that smearing cannot account for a measurable amount of deviation from the initial distribution of the clump sizes. From the observer's point of view we note that Vela~X-1 shows no evidence for an accretion disk, which would show up in the soft X-rays as a blackbody component. The spin period evolution also shows no evidence for an accretion disk, as the pulse period is changing erratically rather than with a constant trend,  which would be expected from the continuous transform of angular momentum from an accretion disk \\citep{ghosh79a, deeter89a}. This lack of an accretion disk supports the argument of a short smearing timescale, because the accretion timescale is also short.  \\subsubsection{Shock structures and grinding} We did not find in the literature any discussion covering the distribution of clump masses in these systems. Assuming that the clumps are log-normal distributed is therefore at the moment an ad hoc assumption and not explained by the models. In this section we will describe a possible mechanism to generate log-normal distributions. To do this, we will first take a closer look at the different features in the accretion region.  Simulations by \\citet{blondin90a, blondin91a} and recently by \\citet{mauche08a} have shown that a shock front forms in front of the neutron star, together with an accretion wake and a possible photoionization wake, as the neutron star ploughs with supersonic motion through the stellar wind of the companion star. These features induce additional variability of the X-ray flux over the orbit, an effect we have seen in Sect. \\ref{sec:phasres}. Even though photoabsorption is negligible at the energies under consideration, Thomson scattering is not. The Thomson cross section depends only weakly on the photon energy, and Thomson scattering can scatter a considerable amount of X-rays out of the line of sight. As shown by \\citet{hanke08a}, deep scattering troughs can be measured in \\inte data of Cygnus X-1, which are ascribed to a large, hot clump passing through the line of sight. From the variations of the median of the orbital-phase-resolved histograms we can calculate the necessary electron column to induce these changes and, assuming ISM abundances, the corresponding \\nh value. The values obtained were on the order of $\\nh=\\left(20-120\\right)\\times10^{22}$\\,cm$^{-2}$, depending on the energy band used.  Compared to that, the average \\nh expected from the stellar wind is $\\nh^{\\text{Wind}} \\approx 2.8\\times10^{22}$\\,cm$^{-2}$, assuming a wind density profile following a $\\beta$-law \\citep{castor75a} with $\\beta = 0.8$.  This result is confirmed by measurements of the X-ray spectrum, which also give values in the regime of 1--$30\\times10^{22}$\\,cm$^{-2}$ \\citep[e.g.,][]{kreykenbohm99a}. Assuming that the variations in our data are solely caused by scattering, the column density must increase up to a factor $\\sim$50 compared to the average wind.  In simulations by \\citet{blondin90a, blondin91a} regions with up to 100 times the density of the ambient wind are seen, which would lead to clumps with radii $r\\approx10^{12}$\\,cm and large masses on the order of $10^{24}$\\,g, assuming spherical clumps. The donor star itself has only a radius of $R_\\mathrm{star} = 2.1\\times10^{12}$\\,cm, making those clumps unrealistically large. The observed variations are thus unlikely to be induced by scattering alone.  The powerful structures in the accretion region will not only change the column density, but will also influence the mass distribution itself. All mass has to pass at least one shock wave and will be subject to strong turbulence prior to accretion. \\citet{blondin91a} have shown that Rayleigh-Taylor instabilities will rise inside the bow shock region, producing overdense eddies and vortices, e.g., structures similar to clumps. It is also very probable that the mass distribution will be changed during this shock transition, as clumps can not only be produced, but large clumps of the stellar wind will be broken up into smaller pieces. Breaking up large clumps is a multiplicative process, and thus forms a log-normal distribution, which is long known in geoscience for the distribution of rocks and sand grains \\citep[and references therein]{smith64a}. \\citet{kevlahan09a} have shown that shocks in the interstellar medium can lead to a log-normal density distribution as well.  Shock fronts and turbulence breaking up clumps can transfer any given distribution into a log-normal like distribution. The emerging masses of the clumps are uncertain however and need more and detailed simulations. A clumpy wind, with huge clumps with masses of at least $10^{22}$\\,g, which is subject to shocks and turbulence grinding down these huge clumps to masses of $\\sim$$10^{20}$\\,g, could therefore be a realistic picture of the accretion geometry.  The clump distribution in the wind, however, would be transformed to a log-normal distribution of the accreted mass by means of the accretion processes, which we can see in the X-ray luminosity.  \\subsubsection{Outlook} Our analysis allows us for the first time to investigate the structure of the wind in an energy band above 20\\,keV. This structure has not been investigated so far in this energy band in systems like Vela~X-1, but shows similar results of a strongly structured wind as other investigations \\citep[e.g.,][]{sako99a}. It is remarkable that the brightness distribution of the black hole HMXB Cygnus~X-1 shows a similar behavior. Although these systems produce their radiation through different processes, the emerging distribution is also log-normal, as shown for short time-scales by \\citet{uttley01a} and \\citet{gleissner04a}.  \\citet{poutanen08a} were able to extend this result to longer time-scales. While our results improve our understanding of the physical processes behind the variations of the luminosity and the structure of the accreted mass, they do not allow to infer where the structure originates from, i.e., from clumps in the stellar wind or from effects during the accretion process. Further investigations should include more detailed models of the brightness distribution and extensive magneto-hydrodynamic simulations of the accretion geometry and the influence of the neutron star on the wind. We propose that a strongly structured wind, with clumps of the masses around $10^{20}$\\,g is accreted onto the neutron star, and thereby subject to shocks and turbulence. These disturbances can lead to a log-normal distribution in accreted masses, which translates directly to the observed luminosity distribution.  In addition, a similar analysis should be performed for other systems, providing more information about that class of objects and their possible connection to SFXT. For example, a preliminary look at 4U 1909+07 \\citep[$\\equiv$ X1908+07,][]{wen00a} has shown that the brightness distribution of this system is following a roughly log-normal distribution as well."
    },
    "1005/1005.2239_arXiv.txt": {
        "abstract": "The statistics of dark matter halos is an essential component of precision cosmology. The mass distribution of halos, as specified by the halo mass function, is a key input for several cosmological probes. The sizes of $N$-body simulations are now such that, for the most part, results need no longer be statistics-limited, but are still subject to various systematic uncertainties. Discrepancies in the results of simulation campaigns for the halo mass function remain in excess of statistical uncertainties and of roughly the same size as the error limits set by near-future observations; we investigate and discuss some of the reasons for these differences. Quantifying error sources and compensating for them as appropriate, we carry out a high-statistics study of dark matter halos from 67 $N$-body simulations to investigate the mass function and its evolution for a reference $\\Lambda$CDM cosmology and for a set of $w$CDM cosmologies. For the reference $\\Lambda$CDM cosmology (close to WMAP5), we quantify the breaking of universality in the form of the mass function as a function of redshift, finding an evolution of as much as $10\\%$ away from the universal form between redshifts $z=0$ and $z=2$. For cosmologies very close to this reference we provide a fitting formula to our results for the (evolving) $\\Lambda$CDM mass function over a mass range of $6\\cdot 10^{11}-3\\cdot 10^{15}$~M$_{\\odot}$ to an estimated accuracy of about 2\\%. The set of $w$CDM cosmologies is taken from the Coyote Universe simulation suite. The mass functions from this suite (which includes a $\\Lambda$CDM cosmology and others with $w\\simeq-1$) are described by the fitting formula for the reference $\\Lambda$CDM case at an accuracy level of 10\\%, but with clear systematic deviations. We argue  that, as a consequence, fitting formulae based on a universal form  for the mass function may have limited utility in high precision cosmological applications. ",
        "introduction": "\\label{section:introduction}  The current paradigm for the formation of cosmological structure is based on the gravitational amplification of primordial density fluctuations in an expanding Universe. The nonlinear transformation of dark matter overdensities -- via a hierarchical dynamical process -- into clumpy distributions called halos, and the subsequent infall of baryons leading to the formation of stars and galaxies within these halos, rounds out the present picture of the formation of observed structure. Although there is no precise mathematical definition of a `halo', several operational definitions -- depending on the particular applications of interest -- have been employed in practice.  The spatio-temporal statistics of halos and sub-halos, as well as of their mass distribution (and its evolution), together provide most of the descriptive framework within which fit all of the structure formation-based probes of cosmology. The mass function alone is a very useful probe in determining cosmological parameters. Because large and massive halos form very late, the high-mass tail of the mass function -- the regime of cluster-scale masses -- is exponentially sensitive to dark energy-related parameters~\\citep{haiman01}. Additionally, the redshift evolution of the cluster mass function depends strongly on the cosmological parameters in a way that is complementary to other probes. The cluster mass function can also be used to measure the normalization of primordial density fluctuations, $\\sigma_8$, and search for hints of primordial non-Gaussianity (see, e.g.,~\\citealt{dalal07, oguri09}). On cluster mass scales and smaller, the mass function, both directly and indirectly, plays an important role in halo models of galaxy formation and statistics, as applied to a wide range of redshifts and objects (predictions of bias, early galaxies, groups, quasars, spatial statistics of luminous galaxies, etc.).  An important motivation for the precision determination of the mass function is the existence of several ongoing and upcoming surveys that aim to detect clusters via their optical, X-ray and Sunyaev-Zel'dovich (SZ) effect signatures~\\citep{maxbcg,abbott05,vikhlinin08,act,spt,bartlett08,fang07}. The number of detected clusters from the individual surveys will range from thousands to tens of thousands. To maximally extract cosmological information from these cluster surveys, the mass function must be specified to better than a few percent accuracy for a range of cosmologies. As discussed by~\\cite{cunha09a} and by~\\cite{wu}, the current theoretical uncertainty in the determination of the mass function can lead to a considerable degradation in the constraints on cosmological parameters. The investigations have also pointed to the usefulness of determining the mass function over a wide range of masses, extending down into the group scale.  Because massive halos are very nonlinear and dynamically nontrivial objects, a fully satisfactory first principles approach to determine the structure and statistics of halos does not yet exist. It follows, therefore, that our current theoretical understanding of the halo mass function is somewhat limited. (This situation may be contrasted to that of nonlinear perturbation theory for the matter power spectrum, where independent of whether individual approaches fail or succeed, the actual problem is well defined conceptually and mathematically.) From the analytical standpoint, the only viable approach to the mass function is still that based on the (heuristic) Press-Schechter (PS) excursion set model \\citep{ps80,bond91} and its extensions (see \\citealt{zentner07} for a review). Although this work has been valuable in suggesting functional forms and representations for the mass function and in analyzing such effects as the scaling of finite-volume corrections, it has not independently yielded predictions for the mass function that are anywhere close to the accuracies that are now required. (For a recent critical assessment, see~\\citealt{robertson08}). Moreover, it is hard to imagine how additional dynamics, gas physics, and feedback mechanisms can be modeled within such a framework. Therefore, it appears that a sufficiently accurate prediction for the mass function of halos can only be achieved using high resolution simulations, and modeling a range of physics tailored to specific applications.  Numerical simulations have become a standard tool to determine the halo mass function over the last decade. Several groups have used suites of simulations to calibrate the halo mass function over an increasingly wider range of masses and redshifts~\\citep{J01,evrard02, white02,reed03,warren05,heitmann06,reed06, lukic07,tinker08,millenium2,crocce09}; see~\\cite{J01} for references to previous work. A key aspect of the calibration of the mass function is the use of $\\ln\\sigma^{-1}(M,z)$ as the central variable, instead of the halo mass, $M$. Here $\\sigma^2(M,z)$ is the variance of the linear density field, extrapolated by linear theory to the redshift of interest, $z$, and smoothed by a spherical top-hat filter of radius $R$, which on average encloses a mass $M$ $(R=[3M/4\\pi\\rho_b(z)]^{1/3})$. The associated scaled differential mass function $f(\\sigma,z;X)$~\\citep{J01} is \\begin{equation} f(\\sigma,z;X)=\\frac{M}{\\rho_b}\\frac{dn_X(M,z)}{d\\ln[\\sigma^{-1}(M,z)]}, \\label{univ} \\end{equation} where $X$ labels the cosmological model and particular halo definition. The variable $\\ln\\sigma^{-1}(M,z)$ appears naturally in the PS approach and extensions thereof, presenting a relatively simple form for $f(\\sigma,z;X)$, in fact one with no dependence on cosmological epoch and parameters. \\cite{J01} found that for a certain fixed definition of halo, independent of cosmology, their simulation results covering redshifts from $z=0-4$, and across different cosmologies, could be well fitted by this ``universal'' form of the mass function at accuracies of order 20\\%. Recent work has shown that mass function universality is apparently not valid beyond the $5-10\\%$ level~\\citep{reed06, lukic07,tinker08, cohn07,crocce09,courtin10}.  Efforts to study this issue further quickly encounter a host of complications. Even independent of such significant physics issues as mass-observable mapping and baryonic effects, it turns out that the choice of halo definition and systematic errors in simulations can easily have as large an effect as that being investigated. Thus, despite the major effort expended in numerical determination of the halo mass function, the present situation cannot be considered to be fully satisfactory, as we discuss in Section~\\ref{section:halos}. Among other sources of error, the effects of finite force resolution and finite sampling error must be carefully dealt with in order to obtain a converged result.  Beyond this point, there is a further cautionary note to keep in mind in terms of precision determination of the mass function: most simulation campaigns have focused on a single cosmology at a time. Therefore, even though results are often quoted in the universal form of Equation~(\\ref{univ}) with small statistical errors, in the absence of rigorous testing of the universal ansatz they cannot be directly applied to cosmologies other than those considered specifically (and even in this case, the actual systematic errors have often turned out to be larger than originally estimated).  Motivated by these considerations, it is important to first establish just how accurately various mass functions can be computed and what the systematic errors are in the most fundamental situation -- the gravity-only $N$-body case. Once an accurate mass function for a particular $\\Lambda$CDM case has been established, it is important to consider a range of observationally relevant redshifts and of cosmologies around that reference point (see, e.g., the discussion in \\citealt{heitmann09}), to understand and explore the range of applicability -- and limitations -- of the (almost) universal description described above. Therefore, the major aims of this paper are: (i) to carefully consider the systematic effects due to numerical errors on the mass function and either avoid or correct for them; (ii) based on these results, establish an accurate prediction for the mass function of a reference $\\Lambda$CDM model at $z=0$; and (iii) extend the investigation to a larger redshift range and provide an accurate prediction for more general $w$CDM models (where the dark energy equation of state parameter, $w$, is constant in time, but $w\\neq -1$).  As a first step, the choice of halo definition has to be considered. For the most part, numerical simulations use two different techniques to identify halos: friends-of-friends (FOF) or spherical overdensity (SO). In the FOF method, halos are found by a percolation technique where particles belong to the same halo if they are within a certain distance (the linking length $b$) of each other. The linking length is typically chosen between $b=0.15$ and $b=0.2$, where $b$ is defined with respect to the mean interparticle spacing. The FOF definition of halos approximately traces isodensity contours and connects more directly to the simulated mass distribution; it is often used in cluster SZ studies. However, the choice of linking length is an issue: too large a linking length can connect neighboring overdensities in a possibly unrealistic manner. The SO method measures the mass in spherical shells around the center of the halo (which is usually determined from the potential minimum of the halo or from the most bound particle) until the density in the shells falls below a certain threshold which is given with respect to either critical or background density. Typically, values for $m_{200}$ to $m_{500}$ (or higher in the case of clusters) are measured (with respect to $\\rho_c$). The spherical overdensity method is particularly convenient for providing predictions for certain kinds of observations, e.g. X-ray cluster masses where one is concerned primarily with studying the inner, virialized region of a halo. The major disadvantage of the SO method is the crudeness of the spherical approximation and that neighboring halos can overlap.  Because isolated, relaxed halos are well-fit by the Navarro-Frenk-White (NFW) profile~\\citep{nfw}, SO and FOF masses are strongly correlated~\\citep{white01, tinker08}. In fact, if the halo concentrations are known, it has been shown that -- in cosmological simulations -- a one-to-one mapping for the two halo definitions exists at the 5\\% level of accuracy~\\citep{lukic08}. However, a fair fraction of halos in simulations are irregular. For currently favored cosmologies, 15-20\\% of $b=0.2$ FOF halos have irregular substructure or have two or more major halo components linked together~\\citep{lukic08}. For such irregular halos, not only does the simple mapping between SO and FOF halos fail, it is not obvious just how to define an appropriate halo mass (lower $b$ to what value, or correspondingly, what choice of overdensity criterion to use?).  In the absence of a compelling theoretical motivation, most numerical studies of the mass function have used FOF masses with linking length $b=0.2$ following the convention set by~\\cite{J01} who noted that this definition led to a universal form for the mass function (for a systematic investigation, see \\citealt{white02, tinker08}). While noting its possible deficiencies, we retain this convention here in order to better compare our results with other work.  Our study of the FOF mass function uses a large suite of runs for a single reference $\\Lambda$CDM cosmology (very close to the WMAP5 parameters from \\citealt{wmap5}) and a set of $w$CDM cosmological simulations -- the ``Coyote Universe'' suite named after the supercomputer it was run on -- that include a different $\\Lambda$CDM model and a few others close to $\\Lambda$CDM. The latter set of simulations represents a simple step beyond $\\Lambda$CDM, where the dark energy equation of state parameter is treated in a purely phenomenological context. Allowing for dark energy evolution (as required by quintessence models, for example) opens up a large parameter space that near-future observations are unlikely to be sensitive to. We, therefore, defer this extension to future work.  In order to carefully control errors, we have followed the criteria for starting redshift, and mass and force resolution as presented in~\\cite{lukic07}. These criteria ensure that halos of a certain size and at certain redshifts can be resolved reliably. In~\\cite{heitmann09} similar criteria were laid out to obtain the matter power spectrum at 1\\% accuracy out to scales $k\\sim 1~h$Mpc$^{-1}$. These criteria are also obeyed by the simulations used here. As discussed further in Section~\\ref{section:halos}, the high-mass tail of the mass function is particularly susceptible to systematic errors in the determination of individual halo masses. These systematic errors can arise from the effects of finite force resolution and we study and characterize these effects. Overall, the contribution of various errors in typical cosmological simulations breaks down basically as follows: (i) too low starting redshift, $\\sim 10\\%$ \\citep{lukic07}, (ii) halo sampling errors, $\\sim 5\\%$, (iii) transfer function approximations, $\\sim 5\\%$, (iv) finite volume effects, $\\sim 1\\%$, (v) force resolution effects, $\\sim 1\\%$. Of these, (i) and (iii) are trivially avoidable, and the others can be controlled at least to the percent level.  In this paper, after compensating for finite sampling and force resolution limitations, we present quantitative results for the mass function from the reference $\\Lambda$CDM simulations and for a suite of $w$CDM cosmologies designed explicitly to bracket the currently observationally relevant range of cosmological parameters \\citep{heitmann09}. We provide a fitting formula describing our reference $\\Lambda$CDM simulation data at the 2\\% accuracy level at the current epoch (we also compute the halo model prediction for the large scale halo bias from the mass function fit). We trace the evolution of the mass function between redshifts $z=0-2$, which represents up to a 10\\% breaking of the universal description of the mass function across a representative range of masses.  We then turn to investigating the variation of the mass function as a function of cosmological parameters from the suite of $w$CDM simulations. This simulation suite, while it lacks some of the statistical power of the reference $\\Lambda$CDM runs, provides a good test of the validity of the universal description of the mass function. At $z=0$, we find that universality for different cosmologies holds to no better than at the 10\\% level, with clear systematic deviations (in both directions) from the quasi-universal form fitted to the reference $\\Lambda$CDM runs. Thus, while not adequate for future precision studies (as is the case for all current mass function fits), our analytic form provides a good estimate for the mass function for $w$CDM cosmologies at the accuracy of currently available data.  \\begin{table}[t] \\begin{center} \\caption{\\label{tab:basic} Parameters for the 38 cosmological models} \\hspace{0cm} \\begin{tabular}{cccccccc} \\tableline\\tableline \\#& $\\omega_m$ & $\\omega_b$ & $n_s$ & $-w$ & $\\sigma_8$ & $h$ & M\\\\ & & & & & & & $10^{14} M_\\odot$\\\\ \\hline 0 & 0.1296 &  0.0224 & 0.9700 & 1.000 &  0.8000 & 0.7200 & 7.00\\\\ 1 & 0.1539 & 0.0231 & 0.9468 & 0.816 & 0.8161  & 0.5977 & 13.3\\\\ 2 & 0.1460 & 0.0227 & 0.8952 & 0.758 & 0.8548 & 0.5970 & 15.9\\\\ 3 & 0.1324 & 0.0235 & 0.9984 & 0.874 & 0.8484 & 0.6763 & 9.96\\\\ 4 & 0.1381 & 0.0227 & 0.9339 & 1.087 & 0.7000 & 0.7204 & 4.42\\\\ 5 & 0.1358 & 0.0216 & 0.9726 & 1.242 & 0.8226 & 0.7669 & 7.20\\\\ 6 & 0.1516 & 0.0229 & 0.9145 & 1.223 & 0.6705 & 0.7040 & 4.27\\\\ 7 & 0.1268 & 0.0223 & 0.9210 & 0.700 & 0.7474 & 0.6189 & 7.30\\\\ 8 & 0.1448 & 0.0223 & 0.9855 & 1.203 & 0.8090 & 0.7218 & 8.04\\\\ 9 & 0.1392 & 0.0234 & 0.9790 & 0.739 & 0.6692 & 0.6127 & 4.98\\\\ 10 & 0.1403 & 0.0218 & 0.8565 & 0.990 & 0.7556 & 0.6695 & 7.58\\\\ 11 & 0.1437 & 0.0234 & 0.8823 & 1.126 & 0.7276 & 0.7177 & 5.64\\\\ 12 & 0.1223 & 0.0225 & 1.0048 & 0.971 & 0.6271 & 0.7396 & 2.26\\\\ 13 & 0.1482 & 0.0221 & 0.9597 & 0.855 & 0.6508 & 0.6107 & 4.78\\\\ 14 & 0.1471 & 0.0233 & 1.0306 & 1.010 & 0.7075 & 0.6688 & 5.42\\\\ 15 & 0.1415 & 0.0230 & 1.0177 & 1.281 & 0.7692 & 0.7737 & 5.47\\\\ 16 & 0.1245 & 0.0218 & 0.9403 & 1.145 & 0.7437 & 0.7929 & 4.22\\\\ 17 & 0.1426 & 0.0215 & 0.9274 & 0.893 & 0.6865 & 0.6305 & 5.50\\\\ 18 & 0.1313 & 0.0216 & 0.8887 & 1.029 & 0.6440 & 0.7136 & 3.05\\\\ 19 & 0.1279 & 0.0232 & 0.8629 & 1.184 & 0.6159 & 0.8120 & 1.88\\\\ 20 & 0.1290 & 0.0220 & 1.0242 & 0.797 & 0.7972 & 0.6442 & 8.24\\\\ 21 & 0.1335 & 0.0221 & 1.0371 & 1.165 & 0.6563 & 0.7601 & 2.80\\\\ 22 & 0.1505 & 0.0225 & 1.0500 & 1.107 & 0.7678 & 0.6736 & 7.46\\\\ 23 & 0.1211 & 0.0220 & 0.9016 & 1.261 & 0.6664 & 0.8694 & 2.19\\\\ 24 & 0.1302 & 0.0226 & 0.9532 & 1.300 & 0.6644 & 0.8380 & 2.44\\\\ 25 & 0.1494 & 0.0217 & 1.0113 & 0.719 & 0.7398 & 0.5724 & 9.09\\\\ 26 & 0.1347 & 0.0232 & 0.9081 & 0.952 & 0.7995 & 0.6931 & 8.24\\\\ 27 & 0.1369 & 0.0224 & 0.8500 & 0.836 & 0.7111 & 0.6387 & 6.28\\\\ 28 & 0.1527 & 0.0222 & 0.8694 & 0.932 & 0.8068 & 0.6189 & 12.6\\\\ 29 & 0.1256 & 0.0228 & 1.0435 & 0.913 & 0.7087 & 0.7067 & 4.14\\\\ 30 & 0.1234 & 0.0230 & 0.8758 & 0.777 & 0.6739 & 0.6626 & 4.09\\\\ 31 & 0.1550 & 0.0219 & 0.9919 & 1.068 & 0.7041 & 0.6394 & 6.25\\\\ 32 & 0.1200 & 0.0229 & 0.9661 & 1.048 & 0.7556 & 0.7901 & 4.30\\\\ 33 & 0.1399 & 0.0225 & 1.0407 & 1.147 & 0.8645  & 0.7286 & 9.37\\\\ 34 & 0.1497 & 0.0227 & 0.9239 & 1.000 & 0.8734 & 0.6510 & 14.5\\\\ 35 & 0.1485 & 0.0221 & 0.9604 & 0.853 & 0.8822 & 0.6100 & 16.4\\\\ 36 & 0.1216 & 0.0233 & 0.9387 & 0.706 & 0.8911 & 0.6421 & 12.9\\\\ 37 & 0.1495 & 0.0228 & 1.0233 & 1.294 & 0.9000 & 0.7313 & 11.7\\\\ \\tableline\\tableline \\vspace{-1.5cm}  \\tablecomments{See text for more details and \\cite{heitmann09} for the model\\\\ selection procedure. To obtain good statistics over a wide range of halo masses,\\\\ the reference $\\Lambda$CDM case, model~0, was augmented by a set of additional runs\\\\ (see Table~\\ref{tab:sim_specs}). The rightmost column shows the mass corresponding to $1/\\sigma=1.8$\\\\ for each cosmology at $z=0$.}  \\end{tabular} \\end{center} \\end{table}  Our results demonstrate that the goal of determining the mass function at the percent level of accuracy will require a much more intensive program of simulations in the future, sampling both cosmological and physical modeling parameters (baryonic physics, feedback), along with well-controlled statistical errors. As has been emphasized earlier \\citep{lukic07}, a possible solution is mass function emulation from a large, but finite set of simulations, using techniques that have been shown to be successful in high-dimensional regression problems \\citep{habib07}.  The paper is organized as follows. In Section~\\ref{section:simulation} we describe the simulation suite used in this paper, encompassing 67 high-resolution simulations for 38 different cosmologies. Several overlapping-volume $\\Lambda$CDM simulations are used to understand and control systematic errors. These errors and their ramifications for the accuracy of the mass determination of halos, and how these translate to limiting the accuracy of the mass function itself, are discussed in Section~\\ref{section:halos}. In Section~\\ref{section:massfunction} we present our results for the mass function for the reference $\\Lambda$CDM model at different redshifts and provide a new fitting form for the mass function matching our simulations at the 2\\% level as well as the associated mass function-derived halo bias. We extend our discussion in Section~\\ref{section:mf_wcdm} to the wider set of $w$CDM cosmologies and investigate how well the mass function fit derived for the reference cosmology holds for this broader class of models. We conclude in Section~\\ref{section:disc}. We discuss error control issues and provide relevant details in Appendix~\\ref{app}. ",
        "conclusions": "\\label{section:disc}   In this work, we study the mass function of dark matter halos over a wide range of masses ($6\\cdot 10^{11}-3\\cdot 10^{15}$~M$_{\\odot}$ for a reference $\\Lambda$CDM model) and a redshift range of $z=0-2$ for a large range of $w$CDM cosmologies. A friends-of-friends algorithm with a linking length of $b=0.2$ is used for halo identification. The primary aims are to control numerical errors and gain sufficient statistical power for cosmological parameter estimation and testing of the universality of the mass function. For a reference $\\Lambda$CDM model, we achieve a 2\\% error in determining the mass function. At this level of numerical control, deviations from universality (in both cosmological parameters and redshift evolution) can be studied systematically. Our level of numerical control is approaching the $N$-body baseline level -- necessary but by no means sufficient -- required by next-generation cosmological surveys. The quest for high accuracy in the `$N$-body' mass function has a natural stopping point at the percent level simply because at this point many other physical processes become important (e.g., baryonic effects, \\citealt{stanek09}). Moreover, the connections to observations need to be directly modeled and end up adding their own significant contribution to the overall error.  We use a large number of high resolution simulations for studying the mass function in a single $\\Lambda$CDM cosmology. These simulations are used to establish the error control methodology in order to obtain an accurate mass function. Error sources systematically studied here include effects of finite force resolution, FOF particle sampling bias, and systematics induced by finite-volume effects.  Using the more accurate mass function fit for this reference cosmology, we also rederive the large scale halo bias using the `peak-backgound split' approach. Compared to previous results obtained with the ST fit, the halo bias changes by 10-15\\% between $z=0-2$, and is systematically lower. Unfortunately, instead of improving the agreement with direct numerical measurements of halo bias (computed by taking ratios of correlation functions), the use of an improved mass function only increases the discrepancy. This points to an essential difficulty with the halo modeling approach. We will return to this problem elsewhere~\\citep{lukic10}.  After studying the mass function in detail for the reference $\\Lambda$CDM cosmology, we extend the range of cosmological parameters by considering a suite of $w$CDM simulations. The range of parameters is set by the current constraints on cosmological parameters. The simulation parameters for the runs are given in Tables~\\ref{tab:basic} and \\ref{tab:sim_specs}. We note that the reference mass function fit provides a good description of the $w$CDM results at an accuracy of $\\sim 10\\%$, however, with systematic deviations that point to clear violations of the universality of the mass function, not only in $w$CDM parameter space, but also within the set of $\\Lambda$CDM models (and models very close to $\\Lambda$CDM).   The breaking of universality as a function of different $\\Lambda$CDM parameters is studied in \\cite{J01, tinker08, warren05}. These studies have shown that universality breaks down at the 20\\% level when $\\sigma_8$ and $\\Omega_m$ are varied. Also \\cite{courtin10} have studied how universality breaks down for quintessence cosmology and found similar variation of universality with cosmology as reported here. The universality of the mass function in the presence of early dark energy has been studied in \\cite{grossi_ede}. We find that universality in the mass function holds at the 10\\% level as a function of cosmological parameters (the $w$CDM suite) and redshift ($w$CDM and the reference $\\Lambda$CDM model).  Given the breaking of universality discussed here (see also \\citealt{tinker08}), it is not clear what the utility of fitting formulae for the mass function might be for observations that actually do require theoretical predictions at the percent level of accuracy. An extensive future simulation campaign will be needed to properly come to grips with this problem. Nevertheless, to provide a compact description of our results, we derive a fitting formula that agrees with our reference $\\Lambda$CDM simulation results at 2\\% accuracy over the redshift range of $z=0-2$ and the mass range of $6\\cdot 10^{11}-3\\cdot 10^{15}$~M$_{\\odot}$. The fitting function is a simple modification of the Sheth-Tormen form with one extra shape parameter to improve the fit at $z=0$ and two extra evolution parameters. This form does not lead to divergences if the normalization condition that all mass resides in dark matter halos is imposed. It holds at the 10\\% level of accuracy for a broad class of $w$CDM cosmologies with the redshift evolution taken into account."
    },
    "1005/1005.0021.txt": {
        "abstract": "We report on the study of an intriguing active galaxy that was selected as a potential multiple supermassive black hole merger in the early-type host SDSS J151709.20+335324.7 ($z=0.135$) from a complete search for double-peaked [O III] lines from the SDSS spectroscopic QSO database. Ground-based SDSS imaging reveals two blue structures on either side of the photometric center of the host galaxy, separated from each other by about $5.7$ kpc. From a combination of SDSS fibre and Keck/HIRES long-slit spectroscopy, it is demonstrated that, in addition to these two features, a third distinct structure surrounds the nucleus of the host galaxy. All three structures exhibit highly-ionized line emission with line ratios characteristic of Seyfert II AGN. The analysis of spatially resolved emission line profiles from the HIRES spectrum reveal three distinct kinematic subcomponents, one at rest and the other two moving at -350 kms$^{-1}$ and 500 kms$^{-1}$ with respect to the systemic velocity of the host galaxy. A comparison of imaging and spectral data confirm a strong association between the kinematic components and the spatial knots, which implies a highly disturbed and complex active region in this object. A comparative analysis of the broadband positions, colors, kinematics and spectral properties of the knots in this system lead to two plausible explanations: a.) a multiple-AGN produced due to a massive dry merger, or, b.) a very powerful radio jet-driven outflow. Subsequent VLA radio imaging reveals a clear jet aligned with the emission line gas, confirming the latter explanation. We use the broadband radio measurements to examine the impact of the jet on the ISM of the host galaxy, and find that the\u00ca energy in the radio lobes can heat a significant fraction of the gas to\u00ca the virial temperature. Finally, we discuss tests that may help future surveys distinguish between jet-driven kinematics and true black-hole binaries. \\j1517\\ is a remarkable laboratory for AGN feedback and warrants deeper follow-up study.  In the Appendix, we present high-resolution radio imaging of a second AGN with double-peaked [O III] lines, SDSS J112939.78+605742.6, which shows a sub-arcsecond radio jet. If the double-peaked nature of the narrow lines in radio-loud AGN are generally due to radio jet interactions, we suggest that extended radio structure should be expected in most of such systems.   %The large spatial seperation between the two components supports the %claim that this is indeed a double AGN, rather than a single active %nucleus with a large extended Narrow-Line Region. A energy budget %analysis provide furthur support for this claim. ",
        "introduction": "Active galactic nuclei are characterized by a non-stellar continuum, narrow emission lines (FWHM $\\lesssim 500~\\mathrm{km~s^{-1}}$), and often broad emission lines (FWHM $\\approx 4000~\\mathrm{km~s^{-1}}$).  The narrow lines generally originate in an extended high ionization region, hundreds of parsecs in size, called the Narrow-Line Region (NLR). In rare cases, the narrow lines from an AGN display a double peaked profile \\citep{zhou04}, with separations between line peaks on the scale of a 100 \\kms\\ or greater, implying complex kinematics of the scale of the NLR. Recent searches for such double-peaked lines in large spectroscopic databases \\citep{liu09, smith09, wang09} uncover that they occur in about a percent of spectroscopic AGN.  Many explanations have been proposed for such double-peaked AGN. These include: unusual kinematics  in the NLR, from a bipolar jet or outflow \\citep{whittle04, das05}, a highly anisotropic NLR geometry \\citep{fws98, xu09} or an unresolved binary AGN, separated on NLR scales \\citep[e.g.,][]{smith09}. Indeed, current double-peaked narrow line samples may include examples of all such processes. For example, \\citet{smith09} find that radio-bright AGN are more common in their sample compared to a control sample of non double-peaked SDSS QSOs. A detailed study of examples of double-peaked narrow-line AGN can help unravel ways to discriminate between different explanations. In particular, the phenomenon of merging super-massive black holes (SMBHs) has important implications for galaxy merger statistics and evolution, and narrow-line kinematics have been used to constrain merger rates at $z \\lesssim 1$ \\citep{comerford09}. In addition, the influence of the AGN on the evolution of its host is an important physical input needed to understand galaxy evolution, but is poorly constrained from observations. Understanding the fraction of double-peaked AGN that are truly merging systems or strong AGN outflows is an important step towards defining reliable samples for studies of galaxy merging and AGN feedback. Here, we present a study of an AGN with strong signatures of a double nucleus, which turns out, after careful analysis, to be a powerful jet-driven outflow.  The intriguing nature of \\j1517\\ was noticed during  a systematic search of QSOs from SDSS Data Release 7 (DR7) to identify objects with double peaked emission-line profiles in  [\\ion{O}{3}] and certain other emission lines \\citep{smith09}. Of the approximately $21,000$ AGN with [\\ion{O}{3}] in the SDSS spectral range, about 40\\% had sufficient quality spectra to show double peaks of a typical nature.  Approximately 1\\%\\ of the objects had definite double-peaked narrow lines. Only two candidates, including \\j1517, showed resolved double nuclei in the SDSS images. \\j1517 was chosen for follow-up spectroscopy because its spectral line profile and appearance in the SDSS images made it appear to be a prime candidate for harboring a double SMBH system.  An interesting point to note is that \\j1517 is a Type II AGN, despite being classified as a Type I QSO in the SDSS database. \\citet{smith09} attribute this to a selection bias which causes double-peaked narrow line objects disproportionately to satisfy the SDSS requirement of an emission line wider than $1000$ \\kms\\ for classification as a QSO. For clarity, in this paper, the term 'double-peaked narrow line AGN' refers to AGN with or without broad lines which have two well-defined peaks in their narrow emission lines.  The organization of this paper is as follows: after introducing the observational data (\\S2), we present basic analyses of the images and spectra (\\S3-5) and a study the emission line profiles (\\S6). An ionization analysis (\\S7) and presentation of new VLA data (\\S8) is followed by a discussion of the system and implications for AGN merger studies (\\S9). We assume a flat $\\Lambda$CDM cosmology with H$_0 = 72$ \\kms\\ Mpc$^{-1}$.   %%%%%%%%%%%%%%%%%%%%% FIGURE 1 %%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure}[t] \\figurenum{1} \\label{filterplot} \\centering \\includegraphics[width=\\columnwidth]{sdss_filterplot.eps} \\caption{ The SDSS spectrum of SDSS J1517+3353. Note the strong emission lines and richness of the line spectrum. The SDSS imaging filter responses are over-plotted. In the lower panel, a zoomed version of the spectrum highlights our fit (in red) to the underlying continuum (in black). The quality of the fit is evident, both in overall shape and the depths of several photospheric features. } \\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\subsection{The Nature of SDSS J1517+3353}  Despite its initial selection as a candidate for a multiple merging black hole system, several lines of evidence confirm that the double-peaked lines of \\j1517\\ are due to a strong emission line outflow produced by a jet-ISM interaction. The radio jet and emission line gas are co-aligned and appear to trace each other quite closely. This is characteristic of Radio Galaxies and Seyferts that show jet-ISM interaction signatures \\citep[][and references therein]{whittle04}. The profiles of the kinematically distinct components appear to show a particular mirror symmetry (\\S5.3), consistent with gas accelerated by a bipolar flow or jet. In addition, the emission line spectrum is consistent with ionization by strong shocks with a mean velocity of $400$ \\kms, consistent with the line widths of the emission line sub-components. Finally, the host galaxy does not appear to display any obvious signatures of recent merger activity. The isophotes are quite regular and the galaxy colors outside the emission line region are quite red. The stellar continuum appears to be that of an early-type galaxy with some extinction in the central regions, presumably due to dust near the nucleus. There is no sign of the substantial recent star-formation which may indicate a gas-rich merger of the sort that can host merging accreting SMBHs.  \\subsection{Analysis of the Interaction}  %%%%%%%%%%%%%%%%%%%%% FIGURE 9 %%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure}[t] \\figurenum{9} \\label{vlaoverlay} \\centering \\includegraphics[width=\\columnwidth]{vla_6cm_u_contours.eps} \\caption{ Contours of the SDSS \\emph{u}-band image (in green) plotted over the 5 GHz radio map. Thin white contours bring out the gradient in the core of the radio source. North is up and East is to the left. A small cross and a circle mark the centers, respectively, of the radio core and the host galaxy (from the SDSS \\emph{z}-band image). Note the strong alignment between ionized gas, traced by the \\emph{u}-band, and radio structures, including the bend in the western jet. } \\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Through a combination of spectral line decomposition and high resolution radio imaging, it is possible to now associate the eastern and western radio jets with specific emission line knots. In this way, a comparison can be made between the dynamics of the emission line gas and the energy content of the radio source.  We begin by estimating the mass and kinetic energies of the ionized gas. The extinction corrected Balmer line luminosity is a direct measure of the total amount of ionized gas:  \\begin{equation} M_{em} = 171\\, F_{H\\beta}\\,d^2 \\, n_{em}^{-1} \\qquad \\textrm{g} \\end{equation}  \\noindent where $d$ is the luminosity distance of the galaxy in cm % 637 Mpc = 1.9655e27 cm and $F_{H\\beta}$ is the H$\\beta$ flux. % Reference for formula? Using the measurements in Table 2, the masses of Feature A/B are both estimated to be around $10^{5.9}$ solar masses. Taking the peak velocity and FWHM as a measure of the motion in the gas, we calculate bulk and turbulent kinetic energies. These are combined to yield total kinetic energies in Features A and B of $10^{54.4}$ and $10^{54.1}$ ergs respectively.  The internal energy of the ionized gas is related to its density and temperature, and can be parametrized by the mass of ionized gas:  \\begin{equation} E_{th} = 1.25\\times 10^{8}\\; M_{em}\\; \\mu^{-1} \\; T_{e}  \\qquad \\textrm{ergs} \\end{equation}  \\noindent where $\\mu \\sim 0.6$ is the mean molecular weight of the gas. For the Features A/B, we get $E_{th} \\sim 10^{51.5}$ for both. This is 2.5-3 orders of magnitude lower than the kinetic energy of the gas. If we assume that the emission line luminosity of the gas is $L_{em} \\approx 10 \\times L_{[O III]} \\sim 10^{43.5}$ erg/s, we derive a timescale at which this internal energy is radiated away $\\tau_{th} \\approx E_{th} / L_{em} = 3$ years. It is clear that the thermal energy of the ionized gas has to replenished, either through jet-driven shocks or photoionization from the central AGN.  The ionized gas energies may be compared to that of the radio source. A listing of the measured flux densities in the four VLA bands in given in Table 3. The mean radio spectral index of the extended emission between 5 and 8 GHz is $\\alpha_{R} = -1.5$. Assuming that this synchrotron spectrum extends between 0.01 GHz and 100 GHz, the total radio luminosity ($L_R$) of the jets can be estimated by normalizing the spectrum to the flux of the jets at 5 GHz, which gives $L_{R} = 10^{42}$ erg s$^{-1}$. From the 5 GHz image, we estimate the size of the eastern jet to be about $1\\farcs2 \\times 1\\farcs1$, which, if we assume a cylindrical jet geometry, gives a radio-emitting volume of $V_{jet} \\approx 15.2$ kpc$^{3}$. Combining this with equipartition energy arguments of \\citet{miley81}, and assuming that the radio emitting plasma has a filling factor $\\sim 1$, mean relativisitic pressures, magnetic fields and energies are estimated to be $P_{me} = 10^{-9.4}$ dyne cm$^{-2}$, $B_{me} = 125$ $\\mu$G and $E_{me} = 10^{56.7}$ ergs. The synchrotron lifetime of the radio source $\\tau_{R} \\approx E_{rel}/L_{R} \\sim 16$ million years, comparable or longer than a typical lifetime for Seyfert-like activity, and much smaller than the timescale on which the radio-source imparts energy to the emission line gas through shocks and pressure-driven acceleration (see below). Synchrotron emission is probably not a major loss mechanism for the energy in the radio plasma. In addition, the energy content of the radio source, assuming equipartition, is more than two orders of magnitude higher than the kinetic energy content of the ionized gas. Only a small fraction of the energy in the lobes needs to go into accelerating the ISM of the galaxy, through radio source expansion or ram pressure-driven shocks, to fully account for the dynamics of the emission line material. We envision a scenario where the jet propagates rapidly to kpc scales, at velocities much larger than the ionized gas, sweeping up the ISM of the host galaxy and shocking it as it expands. The denser component of the shocked ISM in direct contact with the jet is visible as the NLR. There may be a potentially large fraction of shocked gas which cannot cool efficiently by line emission and may remain invisible \\citep[e.g.][]{mellema02}. Therefore, the amount of energy from the jet that is lost to line emission is highly unconstrained.  The pressure of the relativistic plasma, acting over the area of the jet, is the principal source  of momentum for the ionized gas. A rough interaction timescale can be derived by equating $P_{me}$ to the measured bulk momentum in the gas:  \\begin{equation} \\tau = \\frac{M_{em}\\; v_{em}}{P_{me}\\; A_{jet}}   \\qquad \\textrm{seconds} \\end{equation}  \\noindent which yields $\\tau \\sim 10^{5}$ years, for $A_{jet} = 5.3$ kpc$^2$ and an average bulk flow velocity of $v_{em} \\approx 400$ \\kms. The timescale for the acceleration of the gas could be significantly smaller if the major part of the acceleration happened at sub-kpc radii in the galaxy, where $P_{me}$ is expected to be higher. However, the gas would then have to coast out to where it is currently observed on kpc scales at or below its measured bulk velocity of $v_{em}$, adding more than $10^{6}$ years to the overall timescale for the interaction. The timescale would also be longer if the accelerating force was smaller than determined by $P_{me}$, or if the accelerated gas contained a substantial neutral component not detected in emission lines, as implied by some studies \\citep[e.g.][]{morganti05}. Therefore, $\\tau$ is probably a lower limit to the actual time that the jet interaction has been happening in this system.  %XXX CROSSING TIME FOR THE EMISSION BLOBs at 400 km/s is about 3 kpc/0.001 c = 10**7 yr. %Is this timescale useful to compare with the other timescales?YYY  An upper limit to the energy flux $F_E$ of the jet then follows:  \\begin{equation} F_{E} \\quad = \\quad E_{me}/\\tau \\quad = \\quad 1.65 \\times 10^{44} \\: \\textrm{erg s}^{-1} \\end{equation}  This flux is about an order of magnitude larger than the emission line luminosity and two orders of magnitude greater than the radio luminosity. If the timescale of the interaction is indeed as short as $\\tau$, we infer that only a small fraction of the energy output of the jet is lost to shock-excited radiation or other radiative mechanisms.  How does the energy in the lobes compare to the binding energy of the host galaxy? We estimate a velocity dispersion of the galaxy from the NaD absorption line, which is strong and well-defined in the SDSS spectrum. A gaussian fit to the line yields a FWHM of 14.76 \\AA\\, which gives a velocity dispersion $\\sigma_* = 274$ \\kms, after a small correction for the SDSS spectrograph resolution.  The amount of mass that can be unbound by the energy content of the radio source (i.e, $E_{me}$) is given by:  \\begin{equation} M_{bind}  \\; \\approx \\; \\frac{E_{me}}{\\sigma_{*}^{2}} \\end{equation}  \\noindent where we have taken the gravitational binding energy of a unit mass of gas in the central regions of the galaxy to be approximately $\\sigma_{*}^2$ ergs gm$^{-1}$, assuming that the underlying mass density traces a singular isothermal sphere.  Combining our estimates of $E_{me}$ and $\\sigma_{*}$ gives $M_{bind} \\sim 10^{8.5}$ M$_{\\odot}$ which should be accurate to $0.5$ dex, given our uncertainties in the mass distribution of the host galaxy and the velocity dispersion. The energy in the lobes can easily unbind the total mass in ionized gas ($\\sim 10^{6} \\textrm{ M}_{\\odot}$). If we assume a gas fraction in the galaxy of $\\sim 10^{-2}$ and taking its stellar mass to be $10^{11} \\textrm{ M}_{\\odot}$, the gas content of the galaxy, in both hot and cold ISM phases, is estimated to be $10^{9} \\textrm{ M}_{\\odot}$. The energy in the radio lobes can unbind a good fraction of all the gas in the galaxy and is certainly capable of heating most of this gas to its virial temperature, as long as it is effectively coupled with the ISM. The synchrotron lifetime of the lobes sets a simple upper limit to the efficiency of thermalization of the lobe energy. If the radio source is to serve as a source of feedback to the hot atmosphere of the host galaxy, then the energy in the non-thermal population of the lobes has to be exchanged with the surrounding gas within about 16 million years, or else most of it will be lost to radiation, which couples inefficiently with the hot gas.  %Finally, assuming that the jet is relativistic with %$\\gamma = (1-v^2/c^2)^{-1/2} \\sim 3$, the energy content of the jet can be estimated using the following %relationship from \\citet{blandford90}:  %\\begin{equation} %Q = 4\\; P_{me} \\; \\gamma^{2}\\; v\\; A_{j} %\\end{equation}  %\\noindent where $v$ is the jet velocity and $A_{j}$ is the physical cross-sectional area of the jet flow. %An upper limit on $A_j$ comes from the average width of the jet on the 5 GHz image ( $1\\farcs1$) %which yield $A_j \\sim 5.3$ kpc$^{2}$. This gives $\\log Q \\lesssim 46.3$ erg s$^{-1}$. %The energy flux of the jet, assuming that it is relativistic, is four orders of magnitude higher than its radiative %power, implying that much of the energy in the jet can be be injected into its surrounding medium.   \\subsection{Strong AGN Outflows: Implications for Binary SMBH Searches}  \\j1517\\ serves as a case study of a powerful AGN outflow that effectively masquerades as a candidate for merging accreting SMBHs. Only through spatially-resolved spectroscopy and radio imaging is its true nature revealed. In the presence of a jet or strong nuclear outflow, it is much more likely that complex line kinematics on kpc scales are the result of gas interacting with the outflow rather than from inspiralling black holes. \\citet{smith09} find that their double-peaked AGN sample includes more FIRST detected radio bright AGN than the parent QSO population, implying that a non-negligible fraction of the split lines probably come from jet interactions. \\citet{whittle92b} points out that clear non-gravitational kinematics are seen even in radio-quiet Seyferts with 20 cm luminosities $ > 10^{22.5}$ W Hz$^{-1}$.  We suggest that a radio prior be applied to double-peaked samples to identify such outflows. A radio galaxy like \\j1517\\ can be detected to $z\\approx1$ in wide-area 20 cm mJy-level surveys like FIRST and NVSS, while radio-quiet sources above the threshold from \\citet{whittle92b} are detectable with 100 $\\mu$Jy pointed radio observations out to $z \\approx 0.3$ and with wide-area surveys to $z \\approx 0.1$. Modestly deep follow-up of double-peaked narrow-line samples with existing radio facilities to separate outflows from true black hole mergers.  About 10\\% of  local optically-selected Seyfert galaxies have total radio luminosities greater than $10^{22}$ W Hz$^{-1}$ \\citep{debruyn78, rush96, ho01}, significantly higher than the fraction of SDSS AGN which show double-peaked profiles. If only one in ten AGN with bright nuclear radio sources have extended jets, outflows can conceivably account for most of, if not the entire, double-peaked population. Clearly, stronger constraints on the radio content of these samples is needed before they can be used to study merger rates and SMBH evolution.  On the other hand, we have demonstrated that double-peaked emission can indicate the existence of powerful, galaxy-wide outflows, where massive AGN feedback is caught in the act. These samples may dentify the best laboratories with which to directly tackle the details of feedback physics, such as the importance of momentum-driven winds, the effects of anisotropic obscuration, and the efficiency of coupling between the output of the AGN and material in the host galaxy."
    },
    "1005/1005.1847_arXiv.txt": {
        "abstract": "H$_2$ formation is an important process in post-shock regions, since H$_2$ is an active participant in the cooling and shielding of the environment. The onset of H$_2$ formation therefore has a strong effect on the temperature and chemical evolution in the post shock regions. We recently developed a model for H$_2$ formation on a graphite surface in warm conditions. The graphite surface acts as a model system for grains containing large areas of polycyclic aromatic hydrocarbon structures. Here this model is used to obtain a new description of the H$_2$ formation rate as a function of gas temperature that can be implemented in molecular shock models. The H$_2$ formation rate is substantially higher at high gas temperatures as compared to the original implementation of this rate in shock models, because of the introduction of H atoms which are chemically bonded to the grain (chemisorption). Since H$_2$ plays such a key role in the cooling, the increased rate is found to have a substantial effect on the predicted line fluxes of an important coolant in dissociative shocks  [O~\\textsc{I}] at 63.2 and 145.5~$\\mu$m. With the new model a better agreement between model and observations is obtained. Since one of the goals of Herschel/PACS will be to observe these lines with higher spatial resolution and sensitivity than the former observations by ISO-LWS, this more accurate model is very timely to help with the interpretation of these future results. ",
        "introduction": "H$_2$ formation is an important process in post-shock regions, since H$_2$ is an active participant in the cooling and shielding of the environment. The onset of H$_2$ formation therefore has a strong effect on the temperature and chemical evolution in the post shock regions. This letter aims to re-evaluate the estimations and assumptions concerning H$_2$ reformation in molecular shock models. We mainly focus on jump-type (J-type) shocks. These can reach velocities of several hundred km s$^{-2}$ causing an increase of the gas temperature $>10^5$~K. This will, in turn, result in the dissociation of molecular hydrogen, in contrast with continuous-type (C-type) shocks where H$_2$ stays intact. The dissociation of H$_2$ leads to the first cooling of the region \\citep{Kwan:1977, London:1977} to roughly 4300~K. Since all molecular material that is responsible for cooling under normal circumstances (and in C-type shocks), is dissociated, O cooling is the dominant cooling mechanism \\citep{Flower:2003}. At some point, H$_2$ starts to reform again on the surfaces of dust grains. Because of its high abundance, H$_2$ cooling is more efficient than O cooling and it cools the gas phase until roughly 100~K when cooling by H$_2$O takes over. O emission is classically used as a shock tracer, since it requires H$_2$ dissociation and high temperatures, representative for J-type shocks \\citep[$E_{\\rm up} \\sim 230-330$ K; \\emph{e.g.},][]{Nisini:2002}. Observations of O emission are typically complemented by H$_2$ observations, which gives more lines with energy levels spread out over $\\sim$ 500--15\\,000 K. The combination of both observables can, in principle, provide a good handle on the physical conditions in shocks. However, accurate predictions and models are needed to translate these observations to physical conditions.  Detailed molecular shock models can help us with this translation. They provide predictions of observables depending on a set of input parameters describing the physical conditions. Currently, most shock models include a number of ``educated'' estimations for parameter values. This work aims at minimising the uncertainty on at least one of these parameters: the reformation rate of H$_2$ on grain surfaces. As explained above, the exact conditions under which H$_2$  reforms is an important component in the evolutionary sequence and therefore in the interpretation of the observations. Often in gas phase models, the efficiency, \\emph{i.e.}, the fraction of H atoms hitting a grain that leaves the surface as H$_2$, is assumed to be unity and independent of temperature and pressure. Some further assumptions are made on the sticking of hydrogen atoms on the grains.  Grains can be of silicate or carbonaceous nature and they can be covered by a mantle of ice. Here we focus on bare grains, since processing of ices in shocks, like inertial sputtering, has probably lead to the evaporation of the ice mantles. The chemical nature of the grain surface has an effect on the H$_2$ formation rate in terms of both sticking and formation efficiency. We recently developed a model for molecular hydrogen formation on a graphite surface in warm conditions, such as shocks, PDRs and regions around AGNs. The graphite surface acts as a model system for grains containing large areas of polycyclic aromatic hydrocarbon (PAH) structures \\citep{Pendleton:2002}. This type of dust  takes up a large fraction of the interstellar dust surface area and its interaction with hydrogen has been studied substantially in laboratory and theoretical studies. Our kinetic Monte Carlo model was tested against four independent experimental observations (H$_2$ abstraction cross section, saturation coverage, surface structures, and desorption behaviour) and a good agreement was obtained \\citep{Cuppen:2008,Gavardi:2009}. The present letter extends this model to shock conditions by determining the H$_2$ formation efficiencies under these conditions. The improved rate is then introduced into a shock code to study its effect on the temperature and chemistry. The line fluxes of the [O~\\textsc{I}] at 63.2 and 145.5~$\\mu$m are found to be strongly affected by the new H$_2$ formation. In a comparative study between model and observations for HH54, \\cite{Giannini:2006} found an overestimation of the 63.2~$\\mu$m emission by at least a factor of 50. Here we will compare observations of HH46  to our new model and a much better agreement is obtained. ",
        "conclusions": "To summarise, we have introduced a new model for the H$_2$ formation rate on carbonaceous surfaces as a function of gas temperature at a constant low grain temperature.  This is obtained by a state-of-the-art model of H absorption on a graphite surface which includes both physisorption and chemisorption of complex structures. At low gas temperatures, physisorbed atoms can stick to the surface whereas chemisorbed atoms do not have enough energy to overcome the barrier for sticking. The H$_2$  formation rate is dominated by recombination of physisorbed atoms in this regime. At high gas temperatures, most incoming atoms either scatter from the surface or become trapped in a chemisorption well. The H$_2$ formation rate is therefore mainly determined by chemisorbed atoms. This new model for the H$_2$ formation rate is used as input for a molecular shock model. The rate is substantially higher at high gas temperatures as compared to the original implementation of this rate in shock models. This increased rate has a substantial effect on the predicted line fluxes of [O~\\textsc{I}] at 63.2 and 145.5~$\\mu$m and a better agreement between model and observations is obtained than with the previous model. One of the goals of Herschel/PACS will be to observe these lines with higher spatial resolution and sensitivity than the former observations by ISO-LWS.  This more accurate model will help with the interpretation of these future results."
    },
    "1005/1005.4186_arXiv.txt": {
        "abstract": "Non-congruence appears to be the most general form of phase transition in cosmic matter and in the laboratory. In terrestrial applications non-congruence means coexistence of phases with different chemical composition in systems consisting of two (or more) chemical elements. It is just the case for all phase transitions in high-temperature chemically reactive mixtures, which are typical for uranium-bearing compounds in many nuclear energy devices, both contemporary and perspective. As for cosmic matter, most of real and hypothetical phase transitions without nuclear reactions, i.e., those in the interiors of giant planets (solar and extrasolar), those in brown dwarfs and other sub-stellar objects, as well as in the outer crust of compact stars, are very plausible candidates for such type of phase transformations. Two exotic phase transitions, the gas-liquid phase transition in dense nuclear matter and the quark-hadron transition occuring in the interior of compact stars as well as in high-energy heavy-ion collisions are under discussion as the most extreme example of hypothetical non-congruence for phase transformations in High Energy Density Matter. ",
        "introduction": "The term \\emph{non-congruent phase transition} (NCPT) denotes the situation of phase coexistence of two (or more) phases with different chemical compositions. This is a rather evident definition for the case of phase transitions (PT) in most of terrestrial applications (see below) and in astrophysical applications, where nuclear transformations, including $\\beta$-decay, are negligible: PTs in planetary science and outer crust of compact stars etc. The nuclear composition in such situations is conserved and there is no problem with the selection of systems, which fulfill the condition of a NCPT. The situation is more complicated under extreme conditions like in the interiors of compact stars and in remnants of supernova explosions, where nuclear transformations are close to equilibrium. The problem of the NCPT relevance is even more complicated in exotic situations with equilibrium hadron decay and quark deconfinement in interiors of strange (hybrid) stars and in the hypothetical quark-hadron (QH) phase transition in ultrarelativistic heavy-ion collisions at RHIC, LHC, FAIR and NICA. Hence the study of non-congruent phase transitions in typical terrestrial applications could be a useful base for understanding the relevance for such type of phase transitions in exotic situations like interiors of compact stars, supernova explosions, and in the hydrodynamic expansion of a fireball formed in heavy-ion collisions. ",
        "conclusions": ""
    },
    "1005/1005.1595_arXiv.txt": {
        "abstract": "{} {We present the first detailed X-ray analysis of three AGN, the Seyfert~1 galaxies UGC~3142 and ESO~140-43, and the Seyfert~2 galaxy ESO~383-18, in order to study the geometry and the physical characteristics of their absorbers.} {High quality XMM-Newton EPIC and RGS data were analysed, as well as Swift/XRT and BAT and INTEGRAL IBIS/ISGRI data, in order to cover the 0.3--110  keV energy range. For ESO~140-43 also XMM-Newton/OM and Swift/UVOT data were used. We studied the variability of the three AGN on a time-scale of seconds using the EPIC/PN light curves, and the long-term time-scale variability of ESO~140-43 using two observations performed 6 months apart by XMM-Newton.} {The spectra of the three Seyfert galaxies present a \"soft excess'' at energies E~$ < 2 \\rm \\,keV$ above a power-law continuum that can be modeled by complex absorption, without any additional emission component. The X-ray sources in UGC~3142 and ESO~383-18 are absorbed by two layers of neutral material, with covering fractions $\\rm \\,f_1 \\simeq 0.92$ and $\\rm \\,f_2 \\simeq 0.57$ for UGC~3142, and $\\rm  \\,f_{1} \\simeq 0.97$ and $\\rm  \\,f_{2} \\simeq 0.86$ for ESO~383-18. While the clumpy absorber could be part of a disc wind or of the broad line region for UGC~3142, in the case of ESO~383-18 a clumpy torus plus Compton thin dust lanes are more likely. The spectra of ESO~140-43 can be well fitted using a power law absorbed by three clumpy ionized absorbers with different covering factors, column densities, and ionization parameters, likely part of a moving clumpy system, which might be a disc wind or the broad line region. The strong spectral and flux variability on a time scale of 6 months seen in ESO~140-43 is likely due to changes in the moving absorbers. The variation of the covering factor of one of the three ionized absorbers could be detected, on a kilo-seconds time-scale, in the EPIC light-curve of ESO~140-43. } {} ",
        "introduction": "Active Galactic Nuclei (AGN) are thought to be powered by accretion onto super massive black holes. In this picture the X-ray emission is produced by UV photons from the innermost edge of the accretion disc, which undergo multiple inverse Compton scatterings by a population of hot electrons located in a coronal region sandwiching the disc (Haardt \\& Maraschi 1991, 1993). The spectrum in the X-rays can be well approximated by a power law with an exponential cut-off at energy $E_C\\simeq kT_e$, where $T_e$ is the temperature of the electrons. Besides the power-law like continuum, two prominent features generated by the reflection of the continuum are observed in the X-ray spectra of AGN: a neutral iron K$\\alpha$ line (at 6.4 keV in the  local reference frame) and a reflection hump which peaks at E~$\\simeq 30$ keV (Magdziarz \\& Zdziarski 1995). Not all these components are clearly observed in all AGN (e.g., Soldi et al. 2005, Beckmann et al. 2004).  An alternative model for the accretion process onto black holes is that of clumpy accretion flows (e.g., Guilbert \\& Rees 1988). Courvoisier \\& T\\\"urler (2005) assume that the different elements (clumps) of the accretion flow have velocities that may differ substantially. As a consequence, collisions between these clumps will appear when the clumps are close to the central object, resulting in emission  (Ishibashi \\& Courvoisier 2009).  At optical energies, AGN are classified according to their emission lines into type 1 (showing broad and narrow emission lines), and type 2 (showing only narrow lines). A similar distinction between the two types can be done in the X-rays. Here, type 2 objects show absorption which supresses the soft X-ray spectrum. The dividing line between unabsorbed and absorbed AGN is often set at a hydrogen column density of $N_{\\rm H} = 10^{22} \\rm \\, cm^{-2}$. X-ray data show that most AGN unabsorbed in X-ray are optical Seyfert~1 type, and most, but not all, AGN that are absorbed belong to the Seyfert~2 group (e.g., Awaki et al. 1991). Antonucci (1993) proposed that the extinguishing region forms a torus around the central region. The presence of an absorbing torus guarantees anisotropic obscuration of the central region, so that the sources viewed face-on are type 1, and those observed edge-on are type 2 (for which, although it is highly extinguished, the broad line region exists). AGN spectra are mainly affected by absorption at energies $\\lesssim 2$ keV, with the effect depending on whether the absorber is cold (neutral or weakly ionized), or warm (highly ionized). The geometry of the absorber also plays a major role. When it does not cover completely the source (partially covering absorber), some continuum leaks out and reaches the observer, resulting in an apparent excess over the continuum at higher energies. The {\\it Einstein} satellite provided the first evidence that the absorbing gas can cover only a fraction of the line of sight towards the AGN (Reichert et al. 1985, Holt et al. 1980). The partial covering scenario can also explain some of the observed X-ray variability, which could be due to rearrangements of the cloud distribution (e.g., Abrassart \\& Czerny 2000).  The first evidence of absorption due to warm material came from studies of the {\\it Einstein} X-ray spectrum of the QSO MR 2251-178 (Halpern 1984, Pan et al. 1990), and from AGN spectra taken with {\\it ROSAT} (Turner et al. 1993). Since then, highly ionized absorbers have been observed in about half of the X-ray spectra of type 1 AGN, both Seyfert 1 (e.g., Reynolds 1997, George et al. 1998) and quasars (Piconcelli et al. 2005), with column densities up to $\\gtrsim 10^{23} \\rm \\,cm^{-2}$, and often consisting of several zones of ionized gas (e.g., Nandra et al. 1993, Pounds et al. 1994, Kaspi et al. 2002, Steenbrugge et al. 2005). In the case of NGC~3783, Netzer et al. (2003) and Krongold et al. (2003) discussed the observed complexity in the context of several phases of absorbing gas having distinct temperatures and ionization states. Evidence of a multi-phase warm absorber was also recently reported by Longinotti et al. (2010). The observed absorption lines are often blue-shifted with respect to the optical emission lines, which implies that they are produced in an outflowing region, with mean velocities covering the range from hundreds to thousands of $\\rm \\,km \\,s^{-1}$ (Kaspi et al. 2002, Krongold et al. 2003, Blustin et al. 2005, McKernan et al. 2007, Blustin et al. 2007). Analyzing NGC~3516, Turner et al. (2005) discovered a high column absorber with a hydrogen column density of $N_{\\rm \\,H}\\sim 10^{23} cm^{-2}$ and a ionization parameter of $\\log \\xi \\sim 2$, covering about 50\\% of the continuum source. Changes in the covering fraction of the absorber could explain the spectral variability observed in this source (Turner et al. 2008). A complete review on the recent developments of absorption models can be found in Turner \\& Miller (2009).  A soft ($E \\lesssim 2$ keV) excess over the power-law component dominant at higher energies has been found in the X-ray spectra of many Seyfert galaxies (Saxton et al. 1993). The origin of the soft excess is still an open issue. In the past the soft excess was often associated with the high-energy tail of the thermal emission of the disc, but it has been recently shown that the temperature of the disc should be constant (0.1--0.2 keV), regardless of the mass and luminosity of the AGN (Gierlinski \\& Done 2004). This result implies that some other mechanism is at work, as the temperature of the disc should depend on both the mass of the black hole and the accretion rate. Three competing models have been brought forward in order to explain the soft X-ray excess: i) an additional Comptonization component (e.g., Dewangan et al. 2007), ii) ionized reflection (e.g. from the disk, e.g. Crummy et al. 2006), iii) complex and/or ionized absorption (e.g., Gierlinski \\& Done 2004). Chevallier et al. (2006) examined these three main possibilities and showed that, although the absorption model can successfully fit the data of several AGN, for many AGN the quality of the data does not allow a firm conclusion. Nevertheless, they favour a reflection model absorbed by a modest relativistic wind. Done \\& Nayakshin (2007) agree that the excess is most probably due to partially ionized material moving at relativistic speeds close to the black hole. The two potential scenarios for this effect would be: i) an accretion disk seen in reflection, ii) a wind above the disk, seen in absorption. They show that the reflection model would require very specific ionization parameters, contrary to the disk wind model. Schurch \\& Done (2008) pointed out that, although the absorption model using partially ionized material can be physically interpreted as a radiatively driven accretion disc wind, these winds would require velocities of $\\sim 0.9 \\rm c$ in order to reproduce the soft X-ray excess. They conclude that if the soft excess is produced by absorption, it seems more likely that the material is clumpy and/or only partially covers the source rather than forming a continuous outflow.  The three AGN discussed here are part of the {\\it INTEGRAL} AGN catalogue (Beckmann et al. 2009). As at the time of compilation of the catalogue no absorption information was available for a number of sources, unpublished {\\it XMM-Newton} and {\\it Swift}/XRT data were used to constrain $N_{\\rm H}$. The three Seyfert galaxies UGC 3142, ESO 140-43, and ESO 383-18 are those sources which showed particularly strong evidence for excess in the soft X-ray spectra which requires more complex modeling of the absorption.   UGC~3142 is a nearby Seyfert~1 galaxy and was first detected in X-ray by {\\it ROSAT} (Boller et al. 1992), and in hard X-ray by {\\it INTEGRAL} (Bassani et al. 2006). ESO~383-18 is a Seyfert 2 galaxy, detected in hard X-ray by {\\it INTEGRAL} (Krivonos et al. 2007). ESO~140-43, also known as Fairall~51, is a Seyfert 1. The first X-ray detection of this AGN was done by {\\it HEAO A-2} (Marshall et al. 1979), and subsequently it was detected by {\\it ROSAT} (Boller et al. 1992). At hard X-rays ESO~140-43 was detected by {\\it INTEGRAL} IBIS/ISGRI (Beckmann et al. 2007). Information on the 3 sources are reported in Table \\ref{info_sources}. We present here the first detailed X-ray spectral analysis of these three Seyfert AGN using {\\it XMM-Newton}, {\\it Swift} and {\\it INTEGRAL} data. The data analysis is described in Section \\ref{data_analysis}. In Section \\ref{spectral_analysis}, we discuss the spectral modelling of the X-ray spectra, and we present the study of the X-ray and optical/UV variability in Section \\ref{X_UV_variability}. In Section \\ref{discussion}, we discuss the results obtained, and we present the conclusions in Section \\ref{conclusions}.  \\begin{table*} \\caption[]{Positions (J2000), redshift and Galactic hydrogen column densities (from Dickey \\& Lockman 1990) of the 3 sources analyzed in this work.} \\label{info_sources} \\begin{center} \\begin{tabular}{llllll} \\hline Object & Type &RA &  DEC & z & Galactic N$_{\\rm \\,H}$\\\\ & & & & &{\\scriptsize [$10^{21}\\rm \\,cm^{-2}$]} \\\\ \\hline UGC~3142  & Seyfert 1 &$\\rm 04^{\\,h}43^{\\,m}46.8^{\\,s}$ & $+28^{\\circ }58'19\"$ & 0.0217 & 1.9 \\\\ ESO~383-18  & Seyfert 2 &$\\rm 13^{\\,h}33^{\\,m}26.3^{\\,s}$ & $-34^{\\circ }00'59\"$ & 0.0124 & 0.42\\\\ ESO~140-43  & Seyfert 1 &$\\rm 18^{\\,h}44^{\\,m}54.0^{\\,s}$ & $-62^{\\circ }21'53\"$ & 0.0141 & 0.73\\\\ \\hline  \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "We performed the first detailed X-ray analysis of the 3 Seyfert galaxies ESO 140-43, ESO~383-18, and UGC 3142 using {\\it XMM-Newton} high resolution data in combination with {\\it INTEGRAL}/IBIS and {\\it Swift} BAT/XRT data. The spectra of the three AGN presented indeed show what looks like a soft excess at energies E $< 2 \\rm \\,keV$, but which we have shown is due to complex absorption, without any additive emission to the power-law continuum.\\newline The 0.3--10 keV spectrum of the 3 Seyfert galaxies is well fit by a power law, absorbed by two partially covering neutral absorber for UGC~3142 and ESO~383-18, and by 3 clumpy ionized absorbers in the case of ESO~140-43. ESO~140-43 also shows evidence for a reflection hump.  The absorbing medium seen in ESO~383-18 could be a clumpy torus, plus a disc wind or Compton thin dust lanes. For UGC~3142 the observed absorbers might be part of a disc wind, or of the BLR. The variability of the flux and of the hardness ratio observed in the light curve of ESO~383-18 is not explainable by variations of the absorbers, and is likely related to changes of the continuum.  The absorbers of ESO~140-43 have different values of column densities and ionization states, according to the current picture of warm absorbers. Strong flux and spectral (accounted for by changes in the 3 warm absorbers) variability was detected between {\\it ROSAT}, {\\it EXOSAT}, and the two {\\it XMM-Newton} observations of ESO~140-43. The spectral variability is likely due to the fact that the clumpy absorbers are moving, either as a part of a disk wind or of the BLR. During the 2005-observation of ESO~140-43 variability of the covering factor of the LIA is detected on a time-scale of kilo-seconds, while on shorter time-scales the observed spectral variability cannot be explained by changes in the absorbers. Thus the effects of absorption on the variability are observable only on time-scales longer than few kilo-seconds, and are probably hidden by changes of the continuum on shorter time-scales.  }  We reported the first detection of the iron $K\\alpha$ line in the spectra of ESO~383-18 and UGC~3142, and confirmed its presence in the spectrum of ESO~140-43. We tentatively associate an excess around 0.6 keV in the spectrum of ESO~383-18 with OVIII Ly$\\alpha$ emission. More sensitive data will be necessary in order to confirm or rule out the presence of this line. The fact that we observe the iron $K\\alpha$ lines, but not any additive sign of reflection (i.e. a reflection hump) for ESO~383-18 and UGC~3142, is probably due to the low significance of IBIS/ISGRI data and the hard X-rays variability, respectively.  In general there does not seem to be a common model explaining the presence of the soft X-ray excess in all the observed cases. Rather, it is likely that several components play a role, like reflection processes, complex absorption, and additional thermal emission. Not all of these components are observed in all objects and it remains to be clarified, whether the source type and inclination angle of the accretion disk has an influence on the type of soft excess observable in the X-ray spetra. It is anyway likely that some of the reported soft X-ray excesses in Seyfert galaxies are in fact caused by clumpy (neutral or ionized) absorbers, which allow a portion of the primary flux emitted by the central X-ray source to reach the observer directly. Thus the existence of the soft excess might be simply related, at least in some cases, to an incorrect absorption model, which over-estimates the absorbed flux below 2 keV.\\newline New broad band X-ray observations of these 3 Seyfert galaxies, e.g. by {\\it Suzaku}, would be fundamental to probe the hypothesis that the absorption is responsible for the month scale variability, and to give better limits to the sizes and locations of the absorbers."
    },
    "1005/1005.4653_arXiv.txt": {
        "abstract": "We model the potentially observable populations of normal and millisecond radio pulsars in the Large and Small Magellanic Clouds (LMC and SMC) where the known population currently stands at 19 normal radio pulsars.  Taking into account the detection thresholds of previous surveys, and assuming optimal period and luminosity distributions based on studies of Galactic pulsars, we estimate there are $(1.79 \\pm 0.20) \\times 10^4$ and $(1.09 \\pm 0.16) \\times 10^4$ normal pulsars in the LMC and SMC respectively.  When we attempt to correct for beaming effects, and the fraction of high-velocity pulsars which escape the clouds, we estimate birth rates in both the LMC and SMC to be comparable and in the range 0.5--1 pulsar per century. Although higher than estimates for the rate of core-collapse supernovae in the clouds, these pulsar birth rates are consistent with historical supernova observations in the past 300 yr. A substantial population of active radio pulsars (of order a few hundred thousand) have escaped the LMC and SMC and populate the local intergalactic medium. For the millisecond pulsar (MSP) population, the lack of any detections from current surveys leads to respective upper limits (at the 95\\% confidence level) of 15,000 for the LMC and 23,000 for the SMC.  Several MSPs could be detected by a currently ongoing survey of the SMC with improved time and frequency resolution using the Parkes multibeam system. Giant-pulse emitting neutron stars could also be seen by this survey. ",
        "introduction": "\\label{sec:intro}  Our closest neighbouring galaxies, the Large and Small Magellanic Clouds (henceforth LMC and SMC), contain 19 known radio pulsars \\citep{mfl+06}. Since both the LMC and SMC represent a different star formation environment compared to our Galaxy and its globular cluster system, it is interesting to use these results to constrain the pulsar content in the clouds. So far, no millisecond pulsars (MSPs) have been found, though this is in part a result of the relatively low sampling rates and frequency resolutions of previous surveys (see Table~\\ref{tb:param}). The aim of this paper is to use Monte Carlo simulations to reproduce the results from previous LMC and SMC surveys, determine the most likely size of the normal pulsar population, place upper limits on the MSP population, and make predictions for future surveys of the clouds.  In Section~\\ref{sec:surveys}, we compare some of the past surveys of the Magellanic Clouds.  Section~\\ref{sec:normal} contains a detailed description of our simulations and how we are able to reproduce results from the past surveys.  From this model, we are able to infer the population of normal pulsars in the LMC and SMC.  New limits from our MSP simulations are presented in Section~\\ref{sec:msp}.   We discuss our results and present predictions for future Magellanic Cloud surveys in Section~\\ref{sec:discussion}, and finally, in Section~\\ref{sec:conclusions}, we draw our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We successfully reproduce the results from three past Magellanic Cloud surveys, giving us an accurate normal pulsar population for both the LMC and SMC.  These populations provide estimates of the number of normal pulsars located in the LMC and SMC of $(1.79 \\pm 0.20) \\times 10^4$ and $(1.09 \\pm 0.16) \\times 10^4$ pulsars respectively.  Taking into account pulsars whose beams do not intersect our line of sight, and the significant fraction of pulsars which escape both the LMC and SMC, we find that the mean radio pulsar birth rate is most likely to be in the range 0.5--1 pulsars per century in either cloud. This is  consistent with historical observations of supernovae in the Magellanic Clouds. We find that a substantial population of radio pulsars will be ejected from the clouds and populate the local intergalactic medium.  Extending our analysis to the MSP population yields upper limits of 15,000 for the LMC and 23,000 for the SMC.  If the current SMC survey at Parkes produces no MSP detections, that upper limit will be reduced to 8,000.  Consequently, we can realistically expect no more than 3 MSPs to be detected in that survey.  Also, if the current SMC survey extends to the LMC, a maximum of 3 MSPs would be detected, but a lack of MSP detections would reduce the MSP upper limit to 5,500.  Our results here, along with the future results of the ongoing BPSR survey will continue to constrain the number of pulsars, both normal and millisecond, found in the Magellanic Clouds.  With improved simulation techniques and more sensitive surveys, we will be able to model the true pulsar population with increased accuracy to allow for further research in this area. A logical extension of our work would be a full dynamical model of the pulsar content in the Magellanic Clouds as has been recently carried out for radio pulsars in our Galaxy \\citep{fk06,rl10}."
    },
    "1005/1005.1289_arXiv.txt": {
        "abstract": "It has long been regarded as difficult if not impossible for a cosmological model to account simultaneously for the galaxy luminosity, mass, and velocity distributions.  We revisit this issue using a modern compilation of observational data along with the best available large-scale cosmological simulation of dark matter. We find that the standard cosmological model, used in conjunction with halo abundance matching (HAM) and simple dynamical corrections, fits -- at least on average -- all basic statistics of galaxies with circular velocities $\\vcirc > 80\\kms$ calculated at a radius of $\\sim$10~kpc. Our primary observational constraint is the luminosity-velocity relation -- which generalizes the Tully-Fisher and Faber-Jackson relations in allowing all types of galaxies to be included, and provides a fundamental benchmark to be reproduced by any theory of galaxy formation. We have compiled data for a variety of galaxies ranging from dwarf irregulars to giant ellipticals.   The data present a clear monotonic luminosity-velocity relation from $\\sim$50\\kms\\ to $\\sim$500\\kms, with a bend below $\\sim$~80\\kms\\ and a systematic offset between late- and early-type galaxies.  For comparison to theory, we employ our new $\\LCDM$ ``Bolshoi'' simulation of dark matter, which has unprecedented mass and force resolution over a large cosmological volume, while using an up-to-date set of cosmological parameters.  We use halo abundance matching to assign rank-ordered galaxy luminosities to the dark matter halos, a procedure that automatically fits the empirical luminosity function and provides a predicted luminosity-velocity relation that can be checked against observations.  The adiabatic contraction of dark matter halos in response to the infall of the baryons is included as an optional model ingredient.  The resulting predictions for the luminosity-velocity relation are in excellent agreement with the available data on both early-type and late-type galaxies for the luminosity range from $M_r=-14$ to $M_r=-22$.  We also compare our predictions for the ``cold'' baryon mass (i.e., stars and cold gas) of galaxies as a function of circular velocity with the available observations, again finding a very good agreement.  The predicted circular velocity function is also in agreement with the galaxy velocity function from 80 to 400 km s$^{-1}$, using the HIPASS survey for late-type and SDSS for early-type galaxies.  However, in accord with other recent results, we find that the dark matter halos with $V_{\\rm circ}<80 \\kms$ are much more abundant than observed galaxies with the same $V_{\\rm circ}$. Finally, we find that the two-point correlation function of bright galaxies in our model matches very well the results from the final data release of the SDSS, especially when a small amount of scatter is included in the HAM prescription. ",
        "introduction": "\\label{sec:intro}  The cosmological constant + cold dark matter ($\\LCDM$) model is the reigning paradigm of structure formation in the universe. The presence of large amounts of dark mass in the surroundings of galaxies and within galaxy clusters has been established firmly using dynamical mass estimates that include spiral galaxy rotation curves, velocity dispersions of galaxies in clusters and x-ray emission measurements of the hot gas in these systems, as well as strong and weak lensing of background galaxies. The $\\LCDM$ model also correctly predicts the details of the temperature and polarization of the cosmic background radiation \\citep{komatsu10}. A few issues remain where the model and the observations are either hard to reconcile or very difficult to compare \\citep{primack09}. Examples of this are the so-called missing satellites problem \\citep{klypin99,moore99,Bullock00,willman04,Maccio10} and the cusp/core nature of the central density profiles of dwarf galaxies \\citep{flores94,moore94,deBlok97,Valenzuela07,Governato10,deBlok10}.  An outstanding challenge for the $\\LCDM$ model that we address here is to reproduce the observed abundance of galaxies as a function of their overall properties such as dynamical mass, luminosity, stellar mass, and morphology, both nearby and at higher redshifts.  A successful cosmological model should produce agreement with various observed galaxy dynamical scaling laws such as the Faber-Jackson \\citep{faber76} and Tully-Fisher \\citep{tully77} relations.  Making theoretical predictions for properties of galaxies that can be tested against observations is difficult.  While dissipationless simulations can provide remarkably accurate predictions of various properties of dark matter halos, they do not yet make secure predictions about what we actually observe -- the distribution and motions of stars and gas.  We need to find a common ground where theoretical predictions can be confronted with observations.  In this paper we use three statistics to compare theory and observations: (a) the luminosity - circular velocity (LV) relation, (b) the baryonic Tully-Fisher relation (BTF), and (c) the circular velocity function (VF).  In all three cases we need to estimate the circular velocity (a metric of dynamical mass) at some distance from the center of each dark matter halo that hosts a visible galaxy. Unfortunately, theory cannot yet make accurate predictions for the central regions of galaxies because of uncertain baryonic astrophysics.  As a compromise, we propose to use the distance of 10~kpc. Measurements of rotational or circular velocities of galaxies either exist for this distance or can be approximated by extrapolations. At the same time, theoretical predictions at 10~kpc are also simplified because they avoid the complications of the central regions of galaxies.  Our LV relation is a close cousin of the Tully-Fisher (TF) relation and, indeed, we will use some observational results used to construct the traditional Tully-Fisher relation. However, there are substantial differences between the TF and the LV relations. The standard Tully-Fisher relation tells us how quickly spiral disks rotate for given luminosity. The rotation velocity is typically measured at 2.2 disk scale lengths \\citep[e.g.,][]{courteau07}, where the ``cold'' baryons (i.e., stars and cold gas) contribute a substantial fraction of the mass. Instead, at the 10~kpc radius used here for the LV relation, the dark matter is the dominant contribution to the mass in all but the largest galaxies. More importantly, the LV relation includes not only spiral galaxies, but all morphological types. Thus, {\\it the LV relation is the relation between the galaxy luminosity and the total mass inside the 10~kpc fiducial radius.}  In order to make theoretical predictions for the LV relation, we need to estimate the luminosity of a galaxy expected to be hosted by a dark matter halo (including those that are substructures of other halos). There are different ways to make those predictions. Cosmological $N$-body+gasdynamics simulations will eventually be an ideal tool for this.  However, simulations are still far from achieving the resolution and physical understanding necessary to correctly model the small scale physics of galaxy formation and evolution.  Early simulations had problems reproducing the TF relation \\citep[e.g.,][]{navarro00}.  \\citet{eke01} could reproduce the slope of the TF relation, but created disks that were too faint by about $0.5$ magnitudes in the $I$-band for any given circular velocity. Recently the situation has improved. For example, \\citet{governato07} produced disk galaxies spanning a decade in mass that seem to fit both the $I$-band TF relation and the baryonic TF relation very well, as well as the observed abundance of Milky Way-type satellites. Most recently, \\citet{guedes11} have produced perhaps the best match yet to a Milky Way-type galaxy in $\\LCDM$ using a high-resolution smooth particle hydrodynamics simulation with a high density threshold for star formation.  Making predictions for a large ensemble of simulated galaxies is yet another challenge.  Semi-analytical models (SAMs) are a way to make some progress in this direction.  These models have the advantage of producing large-number galaxy statistics.  They typically include many free parameters controlling the strength of the various processes that affect the build-up of the stellar population of a galaxy (i.e., cooling, star formation, feedback, starbursts, AGNs, etc.). Unfortunately, these normalizing parameters can be difficult to constrain observationally \\citep[e.g.,][]{somerville99,benson10}.  The models aim to reproduce the observed number distributions of galaxies as a function of observables such as luminosity, stellar mass, cold gas mass, and half-light radius, along with scaling laws such as the Tully-Fisher relation and the metallicity-luminosity relation.  Early SAMs suffered from serious defects.  The models of \\citet{kauffmann93} were normalized using the observed TF relation zero-point, which resulted in a luminosity function with a very steep faint-end. On the other hand, models such as those of \\citet{cole94} were normalized to reproduce the observed ``knee'' in the LF but this resulted in a large offset in the TF relation zero-point. Later models have shown moderate success in reproducing either the luminosity function \\citep{benson03} or the TF relation \\citep{somerville99}, but it has been difficult to match both simultaneously when rotation curves are treated realistically \\citep{cole00}.  \\citet{benson03} used a combination of disk and halo reheating to obtain reasonable agreement with the observed LF except at the faint end, where they still overpredict the number of dwarf galaxies. If the WMAP 5-year cosmology \\citep{komatsu09} were used, their models would also produce too many very bright galaxies. The TF relation they obtain has the correct slope but their disks are too massive at any given luminosity. Most recently, \\citet{benson10} used a sophisticated version of their GALFORM semi-analytic model to obtain sets of parameters that minimize the deviations from twenty one observational datasets including the LFs in several bands and at different redshifts, the TF relation, the average star formation rate as a function of redshift, clustering and metallicities among many others. Not surprisingly, even their best model has difficulty fitting such a large number of simultaneous constraints. In particular, the LF in the $K$ band overpredicts the number of dwarf galaxies by almost an order of magnitude at the faint end, while the LFs at high redshift consistently overpredict the abundance of all galaxies. In addition, the halos they obtain contain too many satellite galaxies, resulting in too strong a galaxy two-point correlation in the one-halo regime. The Tully-Fisher relation of their best fit model also shows a systematic offset of about $20-40\\kms$ towards higher circular velocities for any given luminosity when compared to observations.  Recent high-resolution $N$-body cosmological simulations such as \\citet{MSI,klypin10} have volumes large enough to obtain the mass function of dark matter (DM) halos, but there is no direct way to compare it to observational measurements of the luminosity or stellar mass functions of galaxies. A new technique recently emerged that allows us to bridge the gap between dark matter halos and galaxies. It is commonly referred to as abundance matching \\citep{Kravtsov04,tasitsiomi04,Vale04,conroy06,Conroy09,Guo10,behroozi10}. Halo abundance matching (HAM) resolves the issue of connecting observed galaxies to simulated dark matter (DM) halos by setting a one-to-one correspondence between red-band luminosity and dynamical mass: more luminous galaxies are assigned to more massive halos.  By construction, it reproduces the observed luminosity function. It also reproduces the scale dependence of galaxy clustering over a range of epochs \\citep{conroy06,Guo10}. When abundance matching is used for the observed stellar mass function \\citep{li09}, it gives a reasonably good fit to the lensing results \\citep{Mandelbaum06} on the relation between the stellar mass and the virial mass \\citep{Guo10}. \\citet{Guo10} also tried to reproduce the observed relation between the stellar mass and the circular velocity with partial success: there were deviations in both the shape and the amplitude. At circular velocities $V_c=100-150 \\kms$ the predicted circular velocity was $\\sim 25$\\% lower than the observed one. They argued that the disagreement is likely due to the fact that they did not include the effect of cold baryons. Below we show that this is indeed the case.  The paper is structured in the following way. Section~\\ref{sec:observations} describes in detail the observational samples used to compare with the results of our analysis. Section~\\ref{sec:sim} briefly describes our new Bolshoi simulation \\citep{klypin10} and compares it to other large cosmological simulations. In section~\\ref{sec:halo} we describe some characteristics of dark matter halos.  Section~\\ref{sec:procedure} describes the abundance matching method used to relate observed galaxies to the DM halos in the Bolshoi simulation and explains the procedure used to measure key quantities such as the circular velocity for these model galaxies. Section~\\ref{sec:observations} describes in detail the observational samples used to compare with the results of our analysis.  Section~\\ref{sec:results} shows the LV relation, the baryonic Tully-Fisher relation, the galaxy circular velocity function and the galaxy two-point correlation function obtained using our model and compares them to the observations described in Section~\\ref{sec:observations}. A brief comparison with related results in the literature is given in Section~\\ref{sec:compare}. Section~\\ref{sec:discussion} presents a discussion of our results and Section~\\ref{sec:conclusions} summarizes them. ",
        "conclusions": "\\label{sec:conclusions}  Here is a short summary of our results:  \\begin{itemize}  \\item In combination with previous results, we conclude that the standard $\\LCDM$ model in conjunction with halo abundance matching can simultaneously fit reasonably well the main global statistics of galaxies: the luminosity function, the stellar mass function, the Luminosity-Velocity relation, the baryonic Tully-Fisher relation, the abundance of galaxies with circular velocities $\\vcirc > 80\\kms$, as well as the clustering properties of bright galaxies as a function of luminosity.  \\item There are systematic deviations in the LV relation with S0 and elliptical galaxies located about 1 magnitude below late types in the LV relation. Massive early types contain less baryonic mass than late types at the same circular velocity.  \\item The range of the effect of contraction of the DM halos due to baryon infall brackets the observations. The LV relation shows preference for a model with moderate contraction, as predicted by \\citet{gnedin04}. This result is compatible with the observed velocity function of galaxies.  \\item There seems to be an overabundance of model galaxies by a factor of $\\sim 10$ compared to observed dwarf galaxies with $\\vcirc < 50\\kms$. This is a serious problem for the $\\LCDM$ model: galaxies with these circular velocities cannot be affected much by ``normal'' physical processes (e.g., supernovae feedback or reionization of the Universe) proposed for the solution of the satellite problem at $\\vcirc \\lesssim 30\\kms$. However, the observational results on the abundance of dwarf galaxies still need to be improved.  \\item Including scatter in luminosity at fixed $\\vc$ using our physically motivated scatter model maintains the agreement of the model LV relation with observations.  \\item The correlation function of the model Bolshoi galaxies matches very well the observations of bright galaxies. The agreement improves when implementing our scatter prescription for all but the brightest galaxies, where a better understanding of the uncertainties is necessary to make a fair test of $\\LCDM$ using halo abundance matching.  \\end{itemize}"
    },
    "1005/1005.2844_arXiv.txt": {
        "abstract": "We provide a set of numerical N-body simulations for studying the formation of the outer Milky Ways's stellar halo through accretion events.  After simulating minor mergers of prograde and retrograde orbiting satellite halo with a Dark Matter main halo, we analyze the signal left by satellite stars in the rotation velocity distribution. The aim is to explore the orbital conditions where a retrograde signal in the outer part of the halo can be obtained, in order to give a possible explanation of the observed rotational properties of the Milky Way stellar halo.  Our results show that, for satellites more massive than $\\sim 1/40$ of the main halo, the dynamical friction has a fundamental role in assembling the final velocity distributions resulting from different orbits and that retrograde satellites moving on low inclination orbits deposit more stars in the outer halo regions end therefore can produce the counter-rotating behavior observed in the outer Milky Way halo. ",
        "introduction": "The Galactic halo has long been considered a single component.  However, evidences in the past few decades have indicated that it may be more complex \\citep{Pre91,Maj92,Kin94,Car96,Chiba00,Kin07,Miceli08}. Recently \\cite{Carollo07} confirmed the existence of a two-components halo, analyzing a large sample of calibration stars from the Sloan Digital Sky Survey (SDSS) DR5 \\citep{Ade07}. According to Carollo et al. the Galactic halo consists of two overlapping structural components, an inner and an outer halo.  These components exhibits different spatial density profiles, stellar orbits and stellar metalicities. In particular the inner-halo stars show a small (or zero) net prograde rotation around the center of the Galaxy.  Outer halo stars possess a clear retrograde net rotation.  The theory of galaxy formation in a Lambda cold dark matter (LCDM) Universe predicts galactic stellar halos to be built from multiple accretion events starting from the first structures to collapse \\citep{White78,Searle78}.  Halos, composed of both dark and baryonic matter, grow by merging with other halos. While the gas from mergers and accretions loses its energy through cooling and settles into a disk, the non dissipative material (accreted stars and dark matter) form a halo around the Galaxy.  In a LCDM universe, a dark matter halo big enough to host the Milky Way contains 300+/-100 subhalos \\citep{Diemand04} \\citep[such a large number of subhalos is related to the so-called \"missing satellites problem\",][]{Moore99}.  There is significant evidence of past accretions into the Milky Way \\citep{Yanny03,Belo06,Grill09}, and therefore we will accept the scenario that the Milky Way halo has been formed by subhalos accretion and we focus on possible origin of a retrograde (counter-rotating) outer stellar halo, defined as the set of stars orbiting around the Galaxy at large distance from the disk plane, say $z \\ga 15$ kpc: \\citet{Carollo07} found that this distance separates the inner and the outer halo. It is expected that such halo is made of stars stripped from merging or disrupted satellites of low mass (\\citet{Zolotov09}; for a detailed study of accreting events onto a Milky-Way like halo, see also \\citet{BoyKol10}). However, from cosmological simulation, prograde and retrograde mergers are approximately equally likely \\citep{Sales07,Read08}.  In order to determine if, and in what condition, we can obtain a retrograde signal in the stellar distribution in the outer halo, we simulate minor mergers of a satellite halo onto a main halo of Galactic mass, with a mass ratio $M_{\\rm primary}/M_{\\rm satellite} \\sim 40$. We put the satellite on two orbits, one prograde and one retrograde. After the satellite completes its merging/disruption process, we identify stars at distance larger than 15 kpc from the disk plane and simply analyze their rotation velocities.  Our main finding is that a counter-rotating stellar halo naturally arises when minor mergers happen on orbits with a low inclination with respect to the disk plane.  The plan of the Letter is the following. In Section \\ref{sec:sim}, we describe our simulations; in Section \\ref{sec:res}, we give our results on counter-rotating outer halo stars, and in Section \\ref{sec:concl} we draw our conclusions. ",
        "conclusions": "\\label{sec:concl}  We performed controlled numerical simulations of minor merger events. Our aim was to determine if an excess of counter-rotating stars in the outer stellar halo can be produced by couples of mergers having orbits with identical inclination and opposite initial rotation. Our primary halo has an exponential stellar disk and its dark matter has a NFW radial density profile; our satellite's DM has an NFW density profile, and contains a stellar spheroid with a Hernquist density profile. The mass ratio of our merger is $M_{\\rm primary}/M_{\\rm satellite} \\sim 40$. We simulate low inclination encounters, in which the angle of the satellite's orbit with the disk plane is 10 degrees, and high inclination ones (60 degrees). We use prograde and retrograde orbits and also vary the DM spin parameter of the main halo. DM spin is aligned with the disk rotation.  Here we define all stars having a distance $z>15$ kpc from the disk plane as belonging to the outer stellar halo.  Our main results are the following: \\begin{itemize}  \\item low inclination mergers do produce an excess of counter-rotating satellite stellar particles in the outer halo, independently on the spin of the DM; \\item our 60 degrees mergers with our 1:40 mass ratio are not able to give a significant counter-rotating signal, also owing to the disk tilt produced by the merger itself; \\item the fraction of counter-rotating to co-rotating satellite stars in the outer stellar halo is higher if the DM has spin.  \\end{itemize}  From our controlled experiments, we now have an indication of the possible origin of the counter-rotation of stars observed in the Galactic outer halo in the context of a hierarchical $\\Lambda$CDM model of galaxy formation, in which many major mergers happen during the history of the Galaxy. Even if, statistically, the number of prograde and retrograde minor merger is the same, still a counter-rotating signal can arise {\\it if such mergers predominantly happen along low inclination orbits}. Since matter accretion, in a CDM dominated Universe, mainly occurs along filaments, this will be the case {\\it if the galaxy disk is co-planar} to the (majority of) filaments. The disk-filament alignment issue is still debated \\cite[see e.g.][]{Brunino07}: from our results, we expect that if the galactic disk were perpendicular to the main accretion streams, no counter-rotating signal should be observed in the outer halo star distribution.  Our orbits are {\\it not} cosmologically motivated, since the aim of our experiment is to determine {\\it if} and {\\it in what cases} an excess of counter-rotating stars in the outer halo can be produced.  Of course, in a realistic case mergers will occur with orbits having a number of different inclinations, giving rise to a velocity distribution which will not show such a clear, double peaked signal as that detected in the present work. A detailed study of the orbital parameters of minor mergers in a statistically significant set of cosmological Galaxy-sized halo cosmological accretion histories is needed before a quantitative comparison between theory expectations and observation can be performed."
    },
    "1005/1005.2419_arXiv.txt": {
        "abstract": "We present a model for the relative velocity of inertial particles in turbulent flows that provides new physical insight into this problem. Our general formulation shows that the relative velocity has contributions from two terms, referred to as the {\\it generalized acceleration} and {\\it generalized shear} terms, because they reduce to the well known acceleration and shear terms in the Saffman-Turner limit. % The generalized shear term represents particles' memory of the flow velocity difference along their trajectories and depends on the inertial particle pair dispersion backward in time. The importance of this backward dispersion in determining the particle relative velocity is emphasized. We find that our model with a two-phase separation behavior, an early ballistic  phase and a later tracer-like phase, as found by recent simulations for the forward (in time) dispersion of inertial particle pairs, gives good fits to the measured relative speeds from simulations at low Reynolds numbers. In the monodisperse case with identical particles, the generalized acceleration term vanishes and the relative velocity is determined by the generalized shear term. At large Reynolds numbers, our model gives a $St^{1/2}$ dependence of the relative velocity on the Stokes number $St$ in the inertial range for both the ballistic behavior and the Richardson separation law. This leads to the same inertial-range scaling for the two-phase separation that well fits the simulation results.  Our calculations for the bidisperse case show that, with the friction timescale of one particle fixed, the relative speed as a function of the other particle's friction time has a dip when the two timescales are similar. This indicates that similar-size particles tend to have stronger velocity correlation than different ones. We find that the primary contribution at the dip, i.e., for similar particles, is from the generalized shear term, while the generalized acceleration term is dominant for particles of very different sizes. Future numerical studies are motivated to check the accuracy of the assumptions made in our model and to investigate the backward-in-time dispersion of inertial particle pairs in turbulent flows. ",
        "introduction": "The dynamics of inertial particles suspended in turbulent flows is of both theoretical and practical importance. Its applications range from industrial processes, e.g., turbulent spray combustion, aerosols and raindrop formation in terrestrial clouds, and dust grain dynamics in astrophysical environments such as interstellar media, protoplanetary disks, and planetary atmospheres. Particle collisions in a turbulent flow are of particular interest because they determine the growth of the particles by coagulation. The estimate of the collision rate requires the understanding of two interesting phenomena regarding inertial particles in turbulence, the preferential concentration and the turbulence-induced relative velocity. The latter is the focus of the present work.  Our primary motivations for this study is its application to theoretical modeling of planetesimal formation in protoplanetory disks. Planetesimals are objects of kilometer size that can further grow into fully-fledged planets.  The formation of these objects is believed to start from the growth of dust grains of micrometer size by collisional coagulation \\cite[e.g.,][]{wei80}. Particles involved in this process have an extensive size range, from micrometer to kilometer,  which corresponds to a range of friction timescale that covers all the scales (from the dissipation range, the inertial range to the outer scales) in the turbulence in protoplanetary disks. Therefore, a detailed understanding of particle collision velocities for a whole range of particle sizes induced by turbulence motions is crucial to investigating particle growth in these disks.  Dust grain collisions in protoplanetary disks do not always lead to coagulation. The grains become less sticky as the size increases. If the relative speed is large, the collision between two large particles may cause them to fragment or to simply bounce off each other \\cite[e.g.,] []{blu08}. To judge the outcome of collisions between large particles thus requires an accurate prediction for the collision speed.  Besides turbulence, there are other effects that can induce relative velocities between particles. For example, gravity gives rise to differential settling for particles of different sizes, which can have important contribution to the relative speed between these particles. In astrophysical environments, radiation pressure and the coupling with magnetic fields (through electric charge on the grain surface) can also play an important role. In the present study, we will neglect these effects, and concentrate on the turbulence-induced relative velocity. We aim at a physical understanding of relative velocities induced by turbulent motions alone, which is clearly a crucial step toward an accurate model for the particle collisions in the presence of both turbulence and the other complexities. The model presented here could be extended to take the other effects into account.  Although our work is motivated by the problem of dust grain collisions in astrophysical environments, it has applications in other contexts such as droplet formation in cloud physics \\cite[e.g.,] []{saf56}.   The relative velocity of two nearby particles induced by turbulent motions has been extensively studied. \\cite{saf56} considered particles with small inertia in the limit that the friction time, $\\tau_p$, of both particles is much smaller than the Kolmogorov timescale, $\\tau_{\\eta}$. % This limit is also expressed as $St \\ll 1$, where the Stokes number, $St$, is defined as the ratio of $\\tau_p$ to $\\tau_{\\eta}$. In this limit, the particle velocity at a given point can be approximately obtained from the (1st-order) Taylor expansion of the particle momentum equation (eq. (2.1) below). For two particles with a separation, $r$, much smaller than the Kolmogorov scale, $\\eta$, \\cite{saf56} derived a result for the average radial relative velocity, $\\langle|w_r|\\rangle$. In the absence of gravity, there are two terms that contribute to the relative speed, a shear term and an acceleration term \\cite[see][]{aya08}, \\begin{equation} \\langle|w_r|\\rangle = \\sqrt{\\frac{2}{\\pi}} \\left (\\frac{1}{15} \\frac {\\bar{\\epsilon}}{\\nu} r^2 + (\\tau_{p2}-\\tau_{p1} )^2 a^2 \\right)^{1/2} \\end{equation} where $\\bar{\\epsilon}$ and $\\nu$ are, respectively, the average dissipation rate and the kinematic viscosity of the turbulent flow, $\\tau_{p1}$ and $\\tau_{p2}$ are the friction timescales of the two particles, and $a$ is the rms acceleration, i.e., $a^2 =\\left \\langle \\left(\\frac{Du}{Dt}\\right)^2 \\right\\rangle$. The factor $\\sqrt{2/\\pi}$ is from the conversion of the radial relative velocity variance, $\\langle w_r^2 \\rangle$, to $\\langle|w_r|\\rangle$ assuming a Gaussian distribution for $w_r$. Note that the shear term is from the longitudinal structure function, $S_{ll}(r)$,  of the flow, which is given by $\\frac{1}{15}\\frac{\\bar{\\epsilon}}{\\nu} r^2 $ for $ r \\lsim \\eta$. In the Saffman-Turner (S-T) limit, the particle velocity follows the local flow velocity very closely, thus the velocity of the two particles at a distance $r<\\eta$ is highly correlated because of the strong flow velocity correlation across a small distance.  The opposite limit is that of heavy particles with friction timescales much larger than the Lagrangian correlation timescale, $T_L$. In this limit, the velocities of two particles coming together are essentially uncorrelated. This is because particles with $\\tau_p \\gg T_L$ have long-time memory, and their current velocities have substantial contributions from the flow velocities on their trajectories in the past. These flow velocities at early times are likely to be uncorrelated because the particles were far away from each other. The relative velocity in this limit is thus determined by the sum of the velocity variances of the two particles, \\begin{equation} \\langle|w_r|\\rangle = \\sqrt{\\frac{2}{\\pi}} \\left [ \\left(v'^{(1)}\\right)^2 + \\left (v'^{(2)} \\right)^2 \\right ]^{1/2} \\end{equation} where $v'^{(1)}$ and $v'^{(2)}$ denote the particle rms velocities. In the limit $\\tau_p \\gg T_L$, they are given by \\cite[e.g.,][]{abr75}, \\begin{equation} \\left(v'^{(1)}\\right)^2  \\simeq u'^2 \\frac{T_L}{\\tau_{p1}} {\\rm ;} \\hspace{1cm} \\left(v'^{(2)}\\right)^2 \\simeq u'^2 \\frac{T_L}{\\tau_{p2}} \\end{equation} where $u'$ is the rms of the flow velocity fluctuations.  In the derivation of eq. (1.3), the temporal correlation of the flow velocity on a particle's trajectory is approximated by the Lagrangian correlation for tracer particles. Discussions on the validity of this assumption will be given in \\S 2.2.  The problem of the relative velocity of inertial particles in these two extreme limits is physically clear, and the results given above are easy to understand and are expected to be generally robust. On the other hand, for particles with intermediate inertia, $\\tau_{\\eta} \\lsim \\tau_p \\lsim T_L$, the problem is more complicated and is less well understood. The velocities of two nearby particles with intermediate $\\tau_p$ are partially correlated and the degree of correlation, intermediate between the two limits, is not easy to evaluate. We will point out that a very important factor in determining this correlation is the distance between the trajectories of the two particles as a function of time before they come close to each other. To our knowledge, this point has not been clearly recognized or explicitly emphasized in the literature. We will show how the separation of two nearby inertial particles backward in time affects the relative velocity between particles with intermediate inertia, $\\tau_{\\eta} \\lsim \\tau_p \\lsim T_L$.  A successful theory for particles of all sizes needs to correctly predict and explain the behavior of the relative velocity between particles with intermediate friction time, as well as recover the results in the two extreme limits. An example of particular theoretical interest is that of identical particles, referred to as the monodisperse case. In the S-T limit ($\\tau_p \\ll \\tau_\\eta$), the acceleration term in eq. (1.1) vanishes for the monodisperse case and the relative velocity does not depend on the friction time. It is constant at a given distance, $r$, and increases linearly with $r$. In the opposite limit with $\\tau_p \\gg T_L$, the relative velocity decreases with the friction time as $\\tau_p^{-1/2}$ or $St^{-1/2}$, according to eqs. (1.2) and (1.3). The question of how $\\langle |w_r| \\rangle$ scales with $\\tau_p$ for $\\tau_p$ in a range corresponding to the inertial range of the turbulent flow, and how it connects with the two extreme limits has not been systematically studied or fully understood. This is one of the primary goals of the present paper. We find that the dispersion of particles backward in time is crucial to answer this question.  The existing models have very different predictions for the relative speed in the inertial range for the monodisperse case \\cite[][]{vol80, wil83, yuu84, kru97, zai03, zai06, aya08}. A detailed discussion of the qualitative differences between these models and their problems will be given in \\S 4.  An important reason for the problems in most of the previous models is that they did not clearly recognize or carefully consider the effect of the particle pair separation backward in time (except for the differential model by \\cite{zai03, zai06} to be discussed below, which we think has the particle backward separation implicitly included). The role of this backward separation will be revealed along the formulation of our model.    In the previous studies that cover a whole range of Stokes numbers, the differential model by Zaichik and collaborators \\cite[][]{zai03a, zai03, zai06} is perhaps the most complete one, as it examines the effect of preferential clustering and the relative speed simultaneously. We will refer to this model as Zaichik {\\it et al}.'s model or the model by Zaichik {\\it et al}. Assuming Gaussian statistics for the flow velocity, the model first sets up an equation for the joint probability distribution function (pdf) of the particle separation and the relative velocity. Deriving the first 3 moment equations of the pdf equation and closing these moment equations by a quasi-normal approximation, Zaichik {\\it et al}. were able to obtain a set of partial differential equations for the particle density correlation function (usually referred to as the radial distribution function) and the particle velocity structure functions. The solution of the differential equations reproduces the two extreme limits discussed above and predicts that the relative velocity of identical particles with intermediate inertia is proportional to $\\tau_p^{1/2}$ or, equivalently, $St^{1/2}$. The validity of this prediction remains to be confirmed by high-resolution simulations. Despite the elegant mathematical formulation and good agreement with results of direct numerical simulations of turbulence with low Reynolds numbers, the model lacks physical transparency in its approximations, especially the quasi-normal assumption. One of the goals of our model is to elucidate the physics behind the inertial range scaling of the relative velocity.    \\cite{fal02} found that an effect, named the sling effect, has a significant contribution to the relative speed \\cite[see also][]{wil05, wil06, fal07}.  The physical picture of the effect is that,  at regions with large negative velocity gradients, faster moving particles can catch up the slower ones from behind \\cite[][]{fal02},  leading to trajectory crossing of the particles (\\cite{bec05}, also see Fig. 1 in \\cite{fal07} for an illustration). This results in a larger relative speed than the S-T prediction for small particles. \\cite{fal07} showed that the effect starts to be important  for $St \\gsim 0.2$ and gives a relative speed several time larger than eq (1.1) for $St$ between $0.2$ and $1$. We will point out a common element shared by the sling effect and our model: the contribution to the relative speed from the particles' memory of the flow velocity difference in the past.  The paper is organized as follows. We present the formulation of our model in \\S 2 (a general formulation in \\S2.1, and basic assumptions and approximations in \\S2.2 and \\S 2.3). The results for identical particles (monodisperse) and different particles (bidisperse) are given in \\S3.1 and \\S 3.2, respectively.  In \\S4, we discuss previous models and compare them with our model. Conclusions are given in \\S 5. ",
        "conclusions": "We have examined the relative velocity of inertial particles suspended in turbulent flows. A general formulation is established based on the calculation of the particle velocity structure function. Our general result for the particle structure function, eq. (2.12), has two terms, a generalized acceleration term, $\\mathcal{A}_{ij}$ (eq. (2.13)), and a generalized shear term, $\\mathcal{D}_{ij}$ (eq. (2.14)). The generalized shear term, $\\mathcal{D}_{ij}$, corresponds to the contribution to the relative speed from particles' memory of  the flow velocity difference in the past.  We find that the backward-in-time dispersion of inertial particle pairs is needed to evaluate this term. The two terms reduce to the acceleration term and the shear term, respectively, in the S-T limit. Our formulation can thus be viewed as a generalization of Saffman and Turner's result for the limit of small particles to particles of any size.      We have shown that our model with a separation behavior similar to that found by recent simulations for the forward (in time) dispersion of inertial pairs \\cite[][]{bec09}, i.e., a ballistic separation followed by a tracer-like behavior, gives quite good fits to the relative speed measured from simulations by \\cite{wan00} (for the monodisperse case) and by \\cite{zho01} (for the bidisperse case).  For the monodisperse case, only $\\mathcal{D}_{ij}$ contributes to the relative velocity. At large Reynolds numbers, a $St^{1/2}$ scaling of the relative velocity in the inertial range is found for both the ballistic separation or the Richardson separation. Therefore, for the two-phase separation that well fits the simulation results, we have the same inertial-range scaling. This scaling is consistent with that from the differenical model by Zaichik and collaborators \\cite[][]{zai03a, zai03, zai06}. Our model provides a clear physical picture for this scaling.   Our calculations for the bidisperse case show that $\\mathcal{A}_{ij}$ dominates the contribution to the relative velocity between particles of very different sizes, while for similar particles the primary contribution is from $\\mathcal{D}_{ij}$. In the relative velocity vs. $St_2$ curves with fixed $St_1$, dips are found around $St_2 \\sim St_1$, indicating stronger velocity correlation for similar size particles than for different size ones. Away from the dips, the relative velocity is essentially given by the contribution from $\\mathcal{A}_{ij}$.  The main assumptions in our model are those for the trajectory correlation and trajectory structure tensors. The approximations for these tensors can be tested and improved by numerical simulations, and our work thus provides a motivation for direct studies of these correlations along the particle trajectories. A direct numerical study of the separation behavior of particles backward in time would also be of interest, because we have shown that it plays an important role  in modeling the relative velocity between particles of similar sizes. With the help of future simulations, the assumptions in our model could be considerably refined. The model may also be extended to include gravity and other effects. The refined and extended model would provide a reliable prediction of the relative velocity between inertial particles and can be applied to many practical studies, such as raindrop formation in atmospheric clouds and collisions of dust grains in astrophysical environments."
    },
    "1005/1005.3734_arXiv.txt": {
        "abstract": "In May 2008, the \\Antares collaboration has completed the construction of the first deep sea neutrino telescope in the Northern hemisphere. \\Antares is a 3D array of 900 photomultipliers held in the sea by twelve mooring lines anchored at a depth of 2500 m in the Mediterranean Sea 40 km off the southern French coast. The detection principle is based on the observation of \\u{C}erenkov light induced by charged particles produced in neutrino interactions in the matter surrounding the detector.  \\vspace{1pc} ",
        "introduction": "Detecting cosmic neutrinos will allow us to obtain information from the inner core of distant sources. Because of their weak coupling to matter, they are not affected neither by massive obstacles nor by background radiation in their way and, as they are not charged, they are not affected by magnetic fields either, and consequently point back to their source. Candidates for neutrino sources are e.g. active galactic nuclei (AGNs) and transient sources like gamma ray bursts (GRB). Hadrons accelerated in these sources interact with the ambient matter and dense photon fields and create pions, which produce neutrinos in their decay. \\begin{figure}[!h] \\includegraphics[width = 17 pc]{antares_scheme_bw.eps} \\label{layout} \\caption{\\Antares detector. A storey is shown in the enlarged view.} \\end{figure} The detection principle of a neutrino telescope such as \\Antares relies on the observation of \\u{C}erenkov light produced by neutrino induced muons by a three dimensional array of photodetectors. The higher the neutrino energy is, the smaller is the angle between the direction of the neutrino and the produced muon.  Cosmic rays interacting with the nuclei in the atmosphere cause a cascade of secondary particles, amongst which high energy muons that constitute an intense source of background in the detector. To suppress this background ($\\sim 1$ per second at the reconstruction level), \\Antares is optimized to detect upward going muons produced by neutrinos within or in the vicinity of the detector after traveling through the Earth. Other sources of background are upward going atmospheric neutrinos ($\\sim 10/$ day), which will have the same signature in the \\Antares detector as the expected cosmic signal. Muons generated by $\\nu_\\mu$ interactions produce a long track in the detector since their interaction length is considerably long in water. ",
        "conclusions": "Since May 2008 the full twelve line \\Antares detector is completed and taking data. First results with the 5 and 10-line configurations are promising as they are comparable with previous experiments, and predictions are optimistic for improved results. The position of \\Antares in the Mediterranean Sea will provide complementary data to those of the Amanda and IceCube detectors, situated in the South Pole.  Different triggers and reconstruction methods are being developed and will be available for the analysis of the twelve-line data."
    },
    "1005/1005.4925_arXiv.txt": {
        "abstract": "Pulsations driven by partial ionization of hydrogen in the envelope are often considered important for driving winds from red supergiants (RSGs). In particular, it has been suggested by some authors that the pulsation growth rate in a RSG can be high enough  to trigger an unusually strong wind (or a super-wind), when the luminosity to mass ratio becomes sufficiently large. Using both hydrostatic and hydrodynamic stellar evolution models with initial masses ranging from 15 to 40~\\Msun, we investigate 1) how the pulsation growth rate depends on the global parameters of supergiant stars, and 2) what would be the consequences of a pulsation-driven super-wind, if it occurred, for the late stages of massive star evolution. We suggest that such a super-wind history would be marked by a runaway increase, followed by a sudden decrease, of the winds mass loss rate.  The impact on the late evolution of massive stars would be substantial, with stars losing a huge fraction of their H-envelope even with a significantly lower initial mass than previously predicted.  This might explain the observed lack of Type II-P supernova progenitors having initial mass higher than about 17~\\Msun.  We also discuss possible implications for a subset of Type IIn supernovae. ",
        "introduction": "Mass loss due to stellar winds is one of the governing factors that determine the evolution of massive stars.  Great improvements in our understanding of line-driven winds from massive hot stars (i.e., O-Type stars or Wolf-Rayet stars) have been achieved over the last three decades, and they are well reflected in recent stellar evolution models \\citep[see][for a recent review]{Puls08}.  However, the driving mechanism for  cool-star winds  remains elusive. Note that many of the authors still rely on the empirical mass-loss rate given by \\citet{deJager88} (hereafter, JNH88), which is more than 20 years old, for modeling cool giant stars, while some alternative prescriptions have been considered by some authors \\citep[e.g.][]{Salasnich99, vanLoon05, vanBeveren07}. Late stages of massive star evolution during and after core helium burning are therefore relatively difficult to study, which is a major obstacle to theoretical investigation on supernova progenitors and their environments.   One of the key factors that should be taken into account for the study of the red supergiant (RSG) evolution is pulsation. It is well known that RSGs are unstable to radial pulsations driven by partial ionization of hydrogen in the outermost layers of the envelope \\citep[e.g.][]{Li94}, and pulsation is one of the most commonly invoked mechanisms for driving winds from cool stars \\citep[cf.][]{Bowen88, Hoefner03, Neilson08}.  Although the detailed mechanisms for driving mass loss by pulsations are not well understood,  simple arguments based on the gain of pulsation energy suggest that RSGs might experience a stronger mass loss for a higher growth rate of the pulsational instability \\citep[e.g.][]{Appenzeller70}. In particular, \\citet{Heger97} found that RSG pulsations can achieve large amplitudes that may be comparable to those in AGB stars when the pulsation periods become close to the thermal time scale of the envelope. They therefore suggested that such strong pulsations might lead to a \"superwind\" if the luminosity to mass ratio in a RSG were sufficiently high. This idea might be supported by observations. For example, \\citet{vanLoon05} found that the dust-enshrouded RSGs in the Large Magellanic Cloud have a very high mass loss rate, which appears to be related to strong RSG pulsation~\\citep{vanLoon08}.    In the present study, we explore possible consequences of such a superwind in the evolution of supernova progenitors. Our approach to this topic is as follows. First, we investigate the properties of RSG pulsations  using our new stellar evolution models that follow the non-linear evolution of pulsation (Sect.~\\ref{sect:properties}).  This allows us to derive a relation between the pulsation growth rate and the structure of RSGs. Second, we assume that the mass loss rate due to RSG winds should be significantly enhanced compared to that of JNH88, when the growth rate of pulsation becomes sufficiently large. The consequent evolution of massive stars can be very different from the standard picture given in previous work (Sect.~\\ref{sect:evol}). Note that here we apply the modified mass loss prescription only to very unstable RSGs, while \\citet{Salasnich99} and \\citet{vanBeveren07} considered overall enhancement of RSG mass loss rate for all types of RSGs. Also, our mass loss prescription is more directly related to the pulsational properties of RSGs than the others that rely on surface luminosity and temperature. Finally, we conclude the paper by discussing implications for supernova progenitors (Sect.~\\ref{sect:discussion}). ",
        "conclusions": "\\label{sect:discussion}  If a pulsationally-driven super-wind (PDSW) phase could be induced by strong pulsations, this would have very important implications for supernova progenitors.  As implied by our model sequences presented above, such a PDSW can significantly reduce the critical ZAMS mass ($M_\\mathrm{crit}$) above which a huge fraction of the stellar envelope is removed before core collapse, thus producing no Type II-P supernova.  Based on our results, we suggest $M_\\mathrm{crit} \\sim$ 19...20\\Msun{} with a PDSW (see Figs.~\\ref{fig:hr} and \\ref{fig:evol}), while the models with the JNH88 mass loss rate predict SN II-P progenitors up to $\\sim$ 25\\Msun. Interestingly, this value of $M_\\mathrm{crit}$ is comparable to what  other alternative  prescriptions of RSG winds mass loss predict \\citep{Salasnich99, vanBeveren07}. Note that $M_\\mathrm{crit}$ could be even further reduced by rotation, as implied by the result of \\citet{Heger97}~\\citep[see also][]{Meynet03}.   This might provide a plausible solution to the so-called RSG problem, i.e. the observed lack of type II-P progenitors with $M_\\mathrm{init} \\ga 16.5 \\pm 1.5$\\Msun \\citep{Smartt09}.  Such observation could also result from the presence of dusty circumstellar material around the stars, as an obscured progenitor would be estimated to have a lower initial mass in pre-SN images.  However \\citet{Smartt09} also noted that the  number of SNe type II-P in their sample is consistent with the expected number of stars in the range 8.5-16.5~\\Msun{} assuming a Salpeter IMF.  This could be a possible indication of the fact that stars with higher initial mass do actually lose most of their hydrogen envelope, which cannot be easily understood with the canonical mass loss rate of JNH88.  The PDSW  is expected to significantly affect the circumstellar medium around a RSG (van Veelen, in prep.).  When the star dies, the shock produced by the collision between the SN ejecta and the circumstellar material transforms kinetic energy into thermal energy. This energy can be radiated away at different wavelengths, resulting in brightnening the supernova remnant for long times \\citep{Chevalier77}.  If the collision occurs directly after the SN explosion, the emission can even alter the spectrum and light curve of the SN. This scenario is usually invoked to explain the sub-class of Type IIn supernovae \\citep{Schlegel90,Filippenko97}.  To reproduce the light curve of the most luminous Type IIn supernovae (SNe IIn) like SN 2006gy and SN 2006tf, very massive shells are required, which need to be ejected in eruptive events a few years before core collapse \\citep[e.g., ][]{vanMarle10}. Pulsational pair instability \\citep{Woosley07} and LBV-like eruptions \\citep[e.g., ][]{Smith07} have been discussed to explain the  presence of shells of 10-25\\Msun{} around the progenitors of luminous SNe IIn.  A PDSW phase is an unlikely explanation for such extreme environments. However it might be interesting for those type IIn where the required mass in the stellar vicinity is of the order of a few solar masses or less.  In this context it is interesting to consider the circumstellar medium around the RSG VY CMa.  The stellar surrounding appears shaped by episodic mass ejections which occurred about 500-1000 yr ago \\citep{Smith09}. The complex morphology of the CSM is suggestive of a possible interaction between convection and pulsation \\citep[see Fig.~13 in][]{Smith09}, which for these stars are predicted to occur on similar timescales \\citep[e.g., ][]{Heger97}. The mass loss rate of $1-2\\times10^{-3}~\\mathrm{M_\\odot yr^{-1}}$ derived by \\citet{Smith09} is comparable to the one expected during the PDSW phase (see Fig.~\\ref{fig:evol}). Even if the mass loss prescription for PDSW we used is somewhat arbitrary, we want to stress that the energy available from the growth of pulsation is enough to drive mass loss rates up to $\\sim 10^{-2}~\\mathrm{M_\\odot yr^{-1}}$. \\citet{Smith09} state that an extreme RSG like VY CMa would produce a Type IIn event like SN 1988Z if it were to explode in its current state. Therefore the occurrence of pulsation-driven super-winds in RSG might explain moderately luminous  SN IIn if the enhanced mass loss takes place less than about $\\sim10^3$ years before core collapse.  With the mass loss prescription of Eq.~\\ref{eq2} with a sufficiently large $\\alpha$ ($\\sim 2$), this would occur in a narrow range around $\\sim$19\\Msun{} for non-rotating stars. An accurate determination of the expected rate of type IIn due to PDSW requires a more realistic mass loss prescription.  To conclude, if a pulsationally-driven super-wind phase could be induced by strong pulsations, we would expect a substantial change in the late evolution of single massive stars. The mass loss rate during the RSG phase would increase dramatically for stars with $M_\\mathrm{init} \\ga M_\\mathrm{crit}$ compared to the JNH88 rate, and the PDSW phase should start earlier in the evolution of more massive stars. The resulting pre-supernova structure of these stars is affected, as well as their CSM.  For single stars, this suggests the following sequence in supernova types as function of increasing initial mass: II-P~$\\xrightarrow{}$~II-L~$\\xrightarrow{}$~IIn~$\\xrightarrow{}$~IIb~$\\xrightarrow{}$~Ib~$\\xrightarrow{}$~Ic. Detailed estimate for the mass range of each supernova requires a better prescription for the pulsationally enhanced mass loss, which deserves further study.  \\vskip 0.5cm We are grateful to Stan Woosley, Bob van Veelen and Norbert Langer for many useful discussions. S.C.Y. is supported by the DOE SciDAC Program (DOE DE-FC02-06ER41438). \\vskip 0.5cm"
    },
    "1005/1005.3028_arXiv.txt": {
        "abstract": "We present the initial imaging and spectroscopic data acquired as part of the VLT VIMOS Lyman-break galaxy Survey. $UBR$ (or $UBVI$) imaging covers five $\\approx36\\arcmin\\times36\\arcmin$ fields centred on bright $z>3$ QSOs, allowing $\\approx21,000$ $2<z<3.5$ galaxy candidates to be selected using the Lyman-break technique. We performed spectroscopic follow-up using VLT VIMOS, measuring redshifts for 1020 $z>2$ Lyman-break galaxies and 10 $z>2$ QSOs from a total of 19 VIMOS pointings. From the galaxy spectra, we observe a $625\\pm510\\kps$ velocity offset between the interstellar absorption and Ly$\\alpha$ emission line redshifts, consistent with previous results. Using the photometric and spectroscopic catalogues, we have analysed the galaxy clustering at $z\\approx3$. The angular correlation function, $w(\\theta)$, is well fit by a double power-law with clustering scale-length, $r_0=3.19^{+0.32}_{-0.54}\\hmpc$ and slope $\\gamma=2.45$ for $r<1\\hmpc$ and $r_0=4.37^{+0.43}_{-0.55}\\hmpc$ with $\\gamma=1.61\\pm0.15$ at larger scales. Using the redshift sample we estimate the semi-projected correlation function, $w_p(\\sigma)$ and, for a $\\gamma=1.8$ power-law, find $r_0=3.67^{+0.23}_{-0.24}\\hmpc$ for the VLT sample and $r_0=3.98^{+0.14}_{-0.15}\\hmpc$ for a combined VLT$+$Keck sample. From $\\xi(s)$ and $\\xi(\\sigma,\\pi)$, and assuming the above $\\xi(r)$ models, we find that the combined VLT and Keck surveys require a galaxy pairwise velocity dispersion of $\\approx700\\kps$, higher than the $\\approx400\\kps$ assumed by previous authors. We also measure a value for the gravitational growth rate parameter of $\\beta(z=3)=0.48\\pm0.17$, again higher than previously found and implying a low value for the bias of $b=2.06^{+1.1}_{-0.5}$. This value is consistent with the galaxy clustering amplitude which gives $b=2.22\\pm0.16$, assuming the standard cosmology, implying that the evolution of the gravitational growth rate is also consistent with Einstein gravity. Finally, we have compared our Lyman-break galaxy  clustering amplitudes with lower redshift measurements and find that the clustering strength is not inconsistent with that of low-redshift $L^*$ spirals for  simple `long-lived' galaxy models. ",
        "introduction": "Observations of the $z\\sim3$ galaxy population present a valuable tool for studying cosmology and galaxy formation and evolution. For cosmology, the interest is in measuring the galaxy clustering amplitudes and redshift space distortions at high redshift. They both lead to virtually independent estimates of the bias whose consistency leads to a test of the standard cosmological model. For theories of galaxy formation and evolution, this is a key period in the history of the Universe in which significant levels of star formation shape both galaxies and the inter-galactic medium (IGM) around them. An especially vital direction of study is the effect of galactic winds at this epoch. Such winds have been directly observed at low \\citep{Heckman90, lehnert99, Martin05, Martin06} and high \\citep{Pettini01, adelberger03, wilman05, adelberger05} redshift and are invoked to explain a range of astrophysical phenomena.  A basic item of cosmological interest is the spatial clustering of the $z\\approx3$ galaxy population itself. In $\\Lambda$CDM, structure in the Universe is known to grow hierarchically through gravitational instability (e.g. \\citealt{mowhite96,1998ApJ...499...20J,2006Natur.440.1137S}) and testing this model requires the measurement of the clustering of matter in the Universe across cosmic time (e.g. \\citealt{2005Natur.435..629S,2008MNRAS.391.1589O,2009MNRAS.400.1527K}). Surveys of matter at $z\\approx3$ currently focus on two main populations, LBGs and Lyman-$\\alpha$ emitters (LAEs). A number of measurements of galaxy clustering are available at $z\\approx3$. For example, \\citet{adelberger03} and \\citet{daangela05b} use the Keck LBG sample with spectroscopic redshfits of \\citet{steidel03} to measure LBG clustering clustering lengths of $r_0=3.96\\pm0.29\\hmpc$ and $r_0=4.48^{+0.17}_{-0.18}\\hmpc$ respectively. Further surveys of LBGs at $z\\approx3$ have produced a range of results with, for example, \\citet{foucaud03} measuring a clustering length for a photometric sample selected from the CFHT Legacy Survey of $r_0=5.9\\pm0.5\\hmpc$, \\citet{2005ApJ...619..697A} measured $r_0 = 4.0 \\pm 0.6 \\hmpc$ at $\\langle z \\rangle = 2.9$ using a different photometric sample whilst \\citet{hildebrandt07} measured a value of $r_0=4.8\\pm0.3\\hmpc$ from an LBG sample taken from GaBoDS data.  \\citet{daangela05b} go on to use the Keck LBG sample to investigate, via redshift space distortions, the gravitational growth rate of the galaxy population at $z\\approx3$, measuring an infall parameter of $\\beta(z=3)=0.25^{+0.05}_{-0.06}$. The infall parameter, $\\beta$, quantifies the large-scale infall towards density inhomogeneities \\citep{hamilton92,hawkins03} and is defined as $\\beta(z)=\\Omega_m(z)^{0.6}/b(z)$, where $\\Omega_m(z)$ is the matter density and $b(z)$ is the bias of the galaxy population. The 2dF Galaxy Redshift Survey (2dFGRS) measurement of the infall parameter nearer the present epoch gave $\\beta(z\\approx0.1)=0.49\\pm0.09$ \\citep{hawkins03}, similar to values obtained by previous local measurements (e.g. \\citealt{1998MNRAS.296..191R}). There have also been dynamical measurements of $\\beta$ at intermediate redshifts using Luminous Red Galaxies where \\citet{ross07} found $\\beta(z=0.55)=0.4\\pm0.05$. \\citet{daangela05a} used the combined 2dF and 2SLAQ QSO redshift surveys to find $\\beta(z=1.5)=0.60\\pm0.14$. Finally, \\citet{2008Natur.451..541G} used the VVDS galaxy redshift survey to measure $\\beta(z=0.77)=0.70\\pm0.26$. As emphasised by \\citet{2008Natur.451..541G}, if there are independent estimates of $b(z)$ for each redshift sample, then the standard model prediction for the evolution with redshift of the gravitational growth rate of $f=\\Omega_m(z)^{0.6}$ can be tested against alternative gravity models.  Here, we shall follow \\citet{daangela05b,daangela05a,2002MNRAS.329..336H} in making their version of the redshift-space distortion cosmological test which also incorporates the \\citet{1979Natur.281..358A} geometric cosmological test.  From redshift-space distortions, we can also determine the small-scale dynamics of the galaxy population which are usually simply modelled as a Gaussian velocity dispersion, measured from the length of the `fingers-of-God' \\citep{1972MNRAS.156P...1J,1987MNRAS.227....1K} in redshift space clustering. This velocity dispersion will generally also include the effects of velocity measurement error. Although \\citet{daangela05b} had to assume a fixed value of $<w_z^2>^{1/2}=400\\kps$ for the mean pairwise velocity dispersion when making their LBG measurement of $\\beta(z=3)$, in bigger surveys it is possble to fit for $<w_z^2>^{1/2}$ and $\\beta$ simultaneously.  Thus in 2dFGRS at $z\\approx0.1$, \\cite{hawkins03} measured a pairwise velocity dispersion of $<w_z^2>^{1/2}\\approx500\\kps$. As well as being of interest cosmologically, the intrinsic galaxy-galaxy velocity dispersion is interesting in terms of establishing the group environment for galaxy formation.  Furthermore, these random peculiar velocities dominate at the smallest spatial scales, significantly affecting clustering measurements on scales $r\\lesssim5\\hmpc$. They influence both the observed galaxy-galaxy clustering and the observed correlation between galaxy positions and nearby Ly$\\alpha$ forest absorption from the IGM (as measured in \\citealt{adelberger03,adelberger05} and \\citealt{2010arXiv1006.4385C}).  To interpret galaxy-IGM clustering results we shall see that measurements of the small scale dynamical velocity dispersion of the galaxy population are very important.  Galactic winds powered by supernovae are a crucial ingredient in models of galaxy formation \\citep{Dekel86, White91}.  Such negative feedback is required to quench the formation of small galaxies and make the observed faint-end of the galaxy luminosity function much flatter than the low-mass end of the dark-matter mass function, see for example the semi-analytical model of \\citet{Cole00}. Simulations without such strong feedback tend to produce galaxies with too massive a bulge, which consequently do not lie on the observed Tully-Fisher relation \\citep{Steinmetz99, Governato10}. Such winds can also remove a significant fraction of baryons from the forming galaxy, thereby explaining why galaxies are missing most of their baryons \\citep{2009astro2010S..25B}, and hence are much fainter in X-ray emission than expected \\citep{2010MNRAS.407.1403C}. In addition, observations of the IGM as probed with QSO sightlines reveal the presence of metals even in the low density regions producing \\lya\\ forest absorption \\citep{songailacowie96, pettini03, aguirre04, 2004A&A...419..811A}. Other than enrichment from galactic scale winds, it is difficult to see from where these metals originate and this is confirmed by simulations (e.g. \\citealt{2009MNRAS.399..574W}).  Direct evidence for outflows in high redshift galaxies came from the Keck LBG survey spectra analysed by \\citet{adelberger03} and \\citet{shapley03} who found evidence for offsets in the positions of ISM absorption lines, Ly$\\alpha$ emission and rest-frame optical emission lines (see also \\citealt{2000ApJ...528...96P,2002ApJ...569..742P}). \\citet{shapley03} present a model in which the optical emission lines arise in nebular star-forming HII regions, giving the intrinsic galaxy redshift, whilst the ISM absorption lines originate from outflowing material surrounding the stellar/nebular component. Ly$\\alpha$ emission arises in the stellar component, but outflowing neutral material scatters and absorbs the blue Ly$\\alpha$ wing, leaving a peak redshifted with respect to the intrisic galaxy redshift \\citep[e.g.][]{2010ApJ...717..289S}. One of our prime aims here is to test the observations underpinning this model in an independent sample of LBGs.  In this paper, we present the first instalment of data of a $z\\sim3$ survey of LBGs within wide ($\\approx30'$) fields centred on bright $z\\sim3$ QSOs. We discuss the imaging and spectroscopic observations, the latter including a search for redshift offsets in the LBG spectra, followed by an analysis of the clustering and dynamics of the LBG galaxy populations in our fields. In a further paper \\citep{2010arXiv1006.4385C}, we present the analysis of the relationship between LBGs and the surrounding IGM via QSO sight-lines, with the intent of further investigating the extent and impact of galactic winds on the IGM.  The structure of this paper is  as follows. We provide the details of our imaging survey in section~\\ref{sec:imaging}, covering observations and data reduction. In section~\\ref{sec:spec}, we present VLT VIMOS spectroscopic observations, describing the data reduction and object identification processes. Section~\\ref{sec:clustering} presents a clustering analysis of the photometrically and spectroscopically identified objects and we finish with our conclusions and summary in section~\\ref{sec:conclusions}. Unless stated otherwise, we use an $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $H_0=100h\\kps\\rm{Mpc}^{-1}$ flat $\\Lambda$CDM cosmology, whilst all magnitudes are quoted in the Vega system. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have described the VLT VIMOS survey of $z\\approx3$ galaxies in a number of fields around bright $z>3$ QSOs. In total this survey has so far produced a total of 1020 LBGs at redshifts of $2<z<3.5$ over a total area of $1.18\\deg2$. This concludes the data acquisition for the initial phase of the VLT VIMOS LBG Survey. At the time of writing, these are the most up to date observations, however the survey has a number of observations only recently acquired, comprising another 25 VIMOS pointings. Upon completion, the survey will comprise a total of 45 VIMOS pointings, building significantly on this initial data-set and providing a catalogue of $\\approx2,000$ $z>2$ galaxies over a sky area of $2.11\\textrm{deg}^{2}$. The wide angular coverage of VLT VIMOS makes the new LBG study very complementary to the previous Keck study which has higher space densities over smaller areas and hence increased power at the smallest LBG separations but little information in the angular direction beyond $10\\hmpc$. We therefore have frequently used the two surveys in combination in the studies of LBG clustering we have presented here.  Based on the fraction of objects observed for this initial VLT LBG survey, we find that our estimated number densities are consistent with previous studies of LBGs in this redshift range. Overall we obtain a mean redshift of $\\overset{\\_}{z}=2.85\\pm0.34$. From the data obtained we have shown evidence for the existence of galactic outflows with comparable offsets between emission and absorption lines as in previous studies (e.g. \\citealt{2000ApJ...528...96P, 2002ApJ...569..742P} and \\citealt{shapley03})  We have further measured  the clustering properties of the VLT VIMOS LBG sample. Based on the angular auto-correlation function of the photometric LBG candidates, the real-space LBG correlation function, $\\xi(r)$, is estimated to take the form of a double power-law, with a break at $r_b\\approx1.5\\hmpc$. This is parametrised by a clustering length and slope below the break of $r_{0,1}=3.19\\pm0.55\\hmpc$, $\\gamma_1=2.45\\pm0.15$ and above the break of $r_{0,2}=4.37^{+0.43}_{-0.55}\\hmpc$, $\\gamma_2=1.61\\pm0.15$.  Assuming $\\gamma=1.8$, the semi-projected  LBG correlation function $w_p(\\sigma)$ gives $r_0=3.67^{+0.23}_{-0.24}\\hmpc$ for the VLT LBGs, slightly lower than $r_0=4.2^{+0.14}_{-0.15}\\hmpc$ for the Keck LBGs, and the combined VLT$+$Keck sample gives $r_0=3.98^{+0.13}_{-0.12}\\hmpc$. At $r_b>1\\hmpc$, the $\\xi(r)$ estimates from $w(\\theta)$ and $w_p(\\sigma)$ are therefore quite consistent. At $r_b<1\\hmpc$ the steeper power-law from the angular correlation function rises above the single power-law that best fits $w_p$, but the difference is only marginally statistically significant. These measurements of LBG clustering are broadly  consistent with previous measurements of the clustering of LBGs at $z\\approx3$ made by \\citet{adelberger03} and \\citet{daangela05b} but  lower than those made by \\citet{foucaud03}  We then  measured the redshift-space LBG auto-correlation function, $\\xi(s)$. As expected, this presents a flatter slope at  scales $s<8\\hmpc$ due to the effect of  velocity errors, outflows and intrinsic velocity dispersions. Both the VLT and Keck samples require total pairwise velocity dispersions in the range $<w_z^2>^{1/2}=600-1000\\kps$ to fit $\\xi(s)$, higher than the $<w_z^2>^{1/2}=400\\kps$ previously assumed \\citep{daangela05b}. The VLT and Keck samples' $\\xi(s)$ results both imply an intrinsic pairwise velocity dispersion of $\\pm400\\kps$ for a $\\xi(r)$ model with $r_0=3.98\\hmpc$ and $\\gamma=1.8$. A higher $<w_z^2>^{1/2}$ will imply a higher  infall parameter, $\\beta(z=3)$, due to the degeneracy between these parameters. The high value of the velocity dispersion will also  have an impact on our search for the effects of star-formation feedback on the QSO Lyman-$\\alpha$ forest \\citep{2010arXiv1006.4385C} because any sharp  decrease in absorption near an LBG will tend to be smoothed away  by this dispersion acting as an effective redshift error.  We combine our LBG sample with that of \\citet{steidel03} with the aim of measuring the infall parameter, $\\beta(z=3)$. Using a  single power-law with $r_0=3.98\\hmpc$ and $\\gamma=1.8$ as our model for the real space $\\xi(r)$, our fits to our measurement of the LBG $\\xi(\\sigma,\\pi)$ from the combined data-set produce a best fitting infall parameter of $\\beta=0.48\\pm0.17$. We find that this value is consistent with the  $\\beta=0.37$ value expected in the standard $\\Lambda$CDM cosmology. For this cosmology the value of the LBG bias implied from the galaxy dynamics is $b=2.06^{+1.12}_{-0.53}$, again consistent with the value of $b=2.22\\pm0.16$ measured from the amplitude of the LBG $\\xi(r)$, assuming the standard cosmology.  We have also made the cosmological test suggested by \\citet{2002MNRAS.329..336H} and \\citet{daangela05b} and shown that the  values of $\\Omega_m^0$ and $\\beta(z=3)$ derived from LBG redshift-space distortion are consistent with those derived by comparing the amplitude of LBG clustering at $z=3$ from the combination of the measured 2dFGRS clustering amplitude and $\\beta$ at $z=0.1$, using linear theory. Our measurement of $\\beta(z=3)$  is therefore consistent with what is expected from the gravitational growth rate predicted by Einstein gravity in the standard cosmological model (see \\citealt{2008Natur.451..541G}).  Finally, we have used the clustering amplitude measured for the LBGs to test simple models of clustering evolution. In particular, we find that if the LBGs are long-lived then they could be the progenitors of low redshift $L^*$ spirals or early-type galaxies by the present day.  The VLT LBG Survey is an ongoing project and we hope to double the survey area and LBG numbers by completion of the project. In combination with this work we are performing a survey of $z\\approx3$ QSOs in our LBG survey fields using the AAOmega instrument at the AAT. Bringing these two data-sets together will present a significant data resource for the study of the relationship between galaxies and the IGM at $z\\approx3$."
    },
    "1005/1005.1394_arXiv.txt": {
        "abstract": "We present 18~GHz observations of the Bullet cluster from the Australia Telescope Compact Array; in particular, a high angular resolution measurement of the substructure in Sunyaev-Zel'dovich Effect (SZE). We report the discovery of a compact SZE `hole' in the galaxy cluster. The SZE hole does not correspond to any bright feature in X-ray, optical or lensing maps; in general, the relatively deeper SZE features appear to avoid the regions with the most intense X-ray emission. These imply that the gas pressure distribution differs significantly from the distributions in gas emission measure, galaxy and dark matter distributions. This has implications for the gas physics and evolution in the cluster merger event. SZE displaced from X-ray centres implies that modeling cluster dynamics is non-trivial; our observations indicate that our current lack of understanding cluster merger astrophysics may be a limitation to modeling the cosmological distribution in SZE cluster counts and the cluster SZE contribution to small-angle cosmic microwave background (CMB) anisotropy. The SZE distribution in the western parts of the cluster are cospatial with the radio halo indicative of a common origin for the hot and relativistic electrons in the turbulent wake of the Bullet. ",
        "introduction": "\\label{intro}  1E0657$-$56, known as the `Bullet Cluster', is one of the hottest known clusters that has been well-studied over the last decade for a variety of reasons; namely, the existence of a strong radio halo \\citep{2000ApJ...544..686L}, the Sunyaev-Zel'dovich effect (\\citet{2009ApJ...701...42H} and references therein), its X-ray brightness \\citep{2002ApJ...567L..27M}, though most notably in providing the most direct proof of the existence of dark matter \\citep{2006ApJ...648L.109C}. It is a cluster collision/merger event at $z\\sim$0.296, with the larger, westward cluster being $\\sim$10 times the mass of the smaller `bullet'.  Observations from the South Pole Telescope \\citep{2009arXiv0912.4317L} show that on arcminute and smaller scales observed CMB anisotropy power is significantly less than levels expected from models of SZE from a cosmological distribution of clusters. This discrepancy may be an indication of our lack of understanding of the environmental impact of cluster mergers. This paper presents sub-arcmin angular resolution observations of the SZE in the Bullet cluster, which demonstrates that the relation between SZE and X-ray emission is complex and requires additional physics that may be related to the impact of the merger on intracluster gas. ",
        "conclusions": "\\label{sz1}  The sky distribution in the SZE does indeed coincide with the radio halo distribution in the western and southen parts; however, SZE is not detected towards the eastern end of the radio halo.  Although the depth of the SZE detection may suffer errors owing to uncertainties in the extrapolation of the halo to 18~GHz, it is clear from our Fig.~\\ref{haloxray} (upper panel) that the detection and extent of the SZE is established.  Assuming an uncertainty of 0.1 in the spectral index of the halo, we infer an uncertainty of $\\pm20$\\% in the peak SZE dip and $5\\arcsec$ in its position. Over most of the SZE distribution as well, the estimated error in SZE decrement is $\\pm20$\\%.  We may derive a temperature scale for the SZE image using the observing frequency and beam FWHM: the image rms noise of 8.4~$\\mu$Jy~beam$^{-1}$ corresponds to 35.3~$\\mu$K. The deepest SZE feature is at RA: $06^{\\rm h}58^{\\rm m}32\\fs3$, DEC: $-55\\degr57\\arcmin03\\arcsec$, and is clearly offset $35\\arcsec$ to the SE of the eastern X-ray emission peak; this deep SZE hole is a flux density decrement of $-$186$\\mu$Jy/beam and corresponds to a brightness temperature decrement of 0.78~mK. The expected SZE temperature decrement given a Compton--$y$ parameter of 3.3$\\times$10$^{-4}$ \\citep{2009ApJ...701...42H} is $\\Delta T_{\\mathrm{SZ}}= -1.983~y~T_{\\mathrm{CMB}}=-1.79$~mK (\\citet{2009ApJ...701...42H} estimate $-$0.77~mK as the peak SZE decrement at 150 GHz from their APEX--SZE observations). Our observed SZE decrement of $-$0.78~mK at 18 GHz is significantly lower.  It may be noted here that our 18~GHz ATCA image of the SZE distribution is an interferometer image and would miss larger scale extended features.  The smallest projected baselines in the ATCA observations correspond to the antenna diameter of 22~m, and the largest angular scale structures that may be reproduced in the ATCA image would be less than the antenna beam FWHM of $2\\farcm6$.  Because of missing large-scale SZE structure the absolute depth of SZE features might be reduced. Significantly differing depths in the SZE decrement in the APEX and ATCA images imply that there is substantial extended SZE structure in the Bullet cluster on angular scales of $2\\farcm6$ and larger, which is resolved out and missing in the ATCA inteferometer image.  In Fig.~\\ref{haloxray} (bottom panel) we show contours of the SZE overlaid on Chandra X-ray intensity distribution shown using grey scales.  \\begin{figure}[!h] \\begin{center} \\includegraphics[height=110mm, angle=0]{45}\\\\ \\caption { {\\it Upper Panel:} Grey scales and contours represent the 18 GHz mosaic image of the SZE in the Bullet cluster; the image has been smoothed to a beam of FWHM $30\\arcsec$. Contours are at ($-$16, $-$12, $-$8, $-$6, $-$4, $-$3, $-$2, 2, 3) times the image rms noise of 8.4~$\\mu$Jy~beam$^{-1}$.  The peaks in ROSAT HRI X-ray emission are marked with cross symbols.\\\\ {\\it Bottom Panel:} Contours represent the 18 GHz mosaic image of the SZE in the Bullet cluster; the image has been smoothed to a beam of FWHM $30\\arcsec$. Contours are at ($-$16, $-$12, $-$8, $-$6, $-$4, $-$3, $-$2, 2, 3) times the image rms noise of 8.4~$\\mu$Jy~beam$^{-1}$. Grey scales represent Chandra X-ray emission. } \\label{haloxray} \\end{center} \\end{figure}  \\begin{enumerate}  \\item We observe a deep SZE `hole' that is significantly displaced to the SE of the X-ray peak in the hotter eastern cluster. Since the SZE probes the integrated pressure along the line of sight, this suggests that there are significant pressure variations in the intracluster gas, presumably because of the ongoing cluster collision. X--ray spectral luminosity is proportional to $\\int n_{\\rm{e}}^{2} T^{-\\frac{1}{2}}d\\rm{l}$, whereas the SZE is proportional to $\\int n_{\\rm{e}}T d\\rm{l}$: the two effects represent different weightings of temperature $T$ and electron density $n_{\\rm{e}}$. X-ray emission images are more representative of $n_{\\rm{e}}$ distribution, whereas the SZE images are representative of $T$ distribution. The offset of the SZE `hole' from the emission peak clearly suggests that the peak in intracluster gas temperature is spatially offset from the peak in gas density in the Bullet cluster\\footnote{\\citep{2002ApJ...567L..27M} have indeed reported finding highest temperatures SE of the X--ray peak.}. Our observations are the first time such an effect has been observed; the high angular resolution ATCA observations of the Bullet cluster presented herein may be the first time that such a detailed comparison between the SZE substructure and X-ray emission has been possible in any cluster. The offset is a demonstration of the complexity in intracluster gas distribution properties in merging clusters. Apart from the deep SZE `hole', we observe significant SZE substructure across the cluster complex. All these SZE components are displaced from the peaks in X-ray emission corresponding to the two clusters comprising the Bullet cluster complex.  \\item All substructure we detect in the SZE image lies within the boundary of the radio halo. Moreover, careful examination of Fig.~\\ref{nsrcyhalo} indicates that the SZE in the western parts of the cluster complex appear to follow the contours of the radio halo.  This is surprising, since the radio halo arises from synchrotron plasma and the SZE from thermal intracluster gas and there is no spatial correspondence expected between these distinct components. A plausible reason for the cospatial locations of the SZE and radio halo---in the western parts of the cluster---is a common origin for the two populations of radiative electrons.  Radio halos are believed to be synchrotron emission from electrons accelerated (or re-accelerated) to relativistic energies in the turbulence generated by cluster mergers.  It may be that the radio halo was created in the turbulent wake of the passage of the `Bullet' through the cluster gas, and the SZE also arises cospatially from relatively hotter electrons in the same turbulent wake. Absence of SZE towards the eastern parts of the radio halo would, in this picture, indicate longer lifetimes for the relativistic electrons compared to the SZE electrons. \\citep{2002ApJ...567L..27M} have reported finding spatial correlation between halo brightness and local gas temperature in merging clusters; this is supplementary to the SZE and the halo being cospatial in our observations.  \\item The SZE decrement is relatively shallow towards the peaks in the X-ray emission.  In Fig.~\\ref{haloxray} it may be noted that the SZE contours are indented in sky regions where the X-ray emission is most intense. The observed avoidance of the highest emission measure regions by SZE suggests that the highest density parts of the cluster gas may have somewhat cooler temperatures, owing to shorter cooling timescales, and consequently contribute relatively less to the SZE. \\end{enumerate} We note here that Massardi et al. (ApJL submitted) found a similar displacement of SZE dip from X--ray brightness peaks in the merging cluster CL0152$-$1357."
    },
    "1005/1005.2959_arXiv.txt": {
        "abstract": "{We present VIMOS integral field spectroscopy of the brightest radio-quiet QSO on the southern sky HE~1029$-$1401 at a redshift of $z=0.086$. Standard decomposition techniques for broad-band imaging are extended to integral field data in order to deblend the QSO and host emission.  We perform a tentative analysis of the stellar continuum finding a young stellar population ($<100$\\,Myr) or a featureless continuum embedded in an old stellar population (10\\,Gyr) typical for a massive elliptical galaxy. The stellar velocity dispersion of $\\sigma_*=320\\pm90$\\,km/s and the estimated black hole mass $\\log(M_\\mathrm{BH}/M_{\\sun})=8.7\\pm0.3$ are consistent with the local $M_\\mathrm{BH}$--$\\sigma_*$ relation within the errors. For the first time we map the two-dimensional ionised gas distribution and the gas velocity field around HE~1029$-$1401. While the stellar host morphology is purely elliptical we find a highly structured distribution of ionised gas out to 16\\,kpc from the QSO. The gas is highly ionised solely by the QSO radiation and has a significantly lower metallicity than would be expected for the stellar mass of the host, indicating an external origin of the gas most likely due to minor mergers. We find a rotating gas disc around the QSO and a dispersion-dominated non-rotating gas component within the central 3\\,kpc. At larger distances the velocity field is heavily disturbed, which could be interpreted as another signature of past minor merger events. Alternatively, the arc-like structure seen in the ionised gas might also be indicative of a large-scale expanding bubble, centred on and possibly driven by the active nucleus.} ",
        "introduction": "Since the underlying host galaxies of several quasi-stellar objects (QSOs) were resolved by \\citet{Kristian:1973}, QSO host galaxies have been extensively studied to understand their properties and their relevance for the evolution of the overall galaxy population over cosmic time. Most of our knowledge of QSO host galaxies comes from numerous ground and space-based imaging studies \\citep[e.g.][]{Malkan:1984,Smith:1986,McLeod:1994,Bahcall:1997,Dunlop:2003,Sanchez:2003}, which provide basic information on the host galaxies like the morphology or the luminosity. Narrow-band images centred on luminous emission lines have been used obtained to infer the distribution and size of the ionised gas surrounding the QSO \\citep[e.g.][]{Stockton:1987,Mulchaey:1996,Falcke:1998,Bennert:2002,Schmitt:2003a}. However, optical and/or near-infrared spectroscopy is required for a more detailed investigation of the properties of the stellar population and the ionised gas.  \\citet{Boroson:1982} were the first to detect strong Balmer absorption lines in the off-nuclear spectrum of the  QSO 3C~48 and highly ionised gas was already found long ago in off-nuclear spectra of many QSOs \\citep[e.g.][]{Wampler:1975,Stockton:1976,Boroson:1984,Boroson:1985} indicating large amount of gas being photoionised  by the QSO radiation. Unfortunately, spectroscopy of QSO hosts suffers invariably from the contamination of the spectrum by the emission of the active galactic nucleus (AGN) even when observed some arcseconds away from the nucleus. Various spectroscopic studies circumvented this problem by looking at obscured (type 2) AGN \\citep[e.g.][]{Gonzalez-Delgado:2001,Kauffmann:2003,CidFernandes:2004,Bennert:2006a,Stoklasova:2009}, which are thought to have the same properties as unobscured (type 1) AGN in the framework of the unification model of AGN \\citep{Antonucci:1993,Urry:1995}. They often find a young stellar population in the AGN host galaxies, in particular for the most luminous AGN, suggesting a connection between star formation and nuclear activity in agreement with photometric studies of type 1 QSOs at all redshifts \\citep{Jahnke:2003,Jahnke:2004b,Jahnke:2004c,Sanchez:2004b}.  \\citet{Letawe:2007} and \\citet{Jahnke:2007}, hereafter Let07 and Jah07 respectively, developed two different methods to decompose the AGN and host spectra in long-slit observations and applied them to a sample of luminous QSOs. However, long-slit spectroscopy covers only a small part of the host on the chosen position angle and slit width. Only integral field unit (IFU) observations allow to study the properties of the entire host galaxy in an unbiased way. Dedicated AGN-host decomposition techniques have already been developed and were successfully applied to IFU data of luminous QSOs at low and high redshift \\citep{Sanchez:2004a,Christensen:2006,Husemann:2008}.  In this paper, we report on optical IFU spectroscopy of the luminous radio-quiet QSO \\object{HE 1029$-$1401} with the VIMOS IFU at the Very Large Telescope (VLT). HE~1029$-$1401 was discovered by the Hamburg-ESO survey \\citep[HES;][]{Wisotzki:1991,Wisotzki:2000} to be the brightest QSOs in the southern hemisphere with an apparent magnitude of $V=13.7$ at a redshift of only $z=0.086$.  Observations with the Hubble Space Telescope revealed a bright elliptical (E1) host galaxy \\citep{Bahcall:1997} with an effective radius of $r_\\mathrm{e}=3.2\\arcsec$ (5.1\\,kpc). \\citet{Jahnke:2004b} found from multi-colour imaging that the host colours are significantly bluer than inactive galaxies at the same luminosity and inferred an intermediate-age stellar population of 0.7\\,Gyr (luminosity weighted) by modelling the spectral Energy distribution (SED) with template Single Stellar Populations (SPPs).  Spatially resolved on-nuclear longslit spectroscopy for this object was presented by Let07 and Jah07. Both studies consistently found highly ionised gas in the host galaxy on kpc scales. Jah07 found that the stellar component is non-rotating, but  the velocity curve of the ionised gas indicated the presence of a rotating gas disc on kpc scales around the nucleus. Let07 argued instead that the velocity curve does not fit with a pure rotational velocity field based on inclination arguments.  The main focus of this paper is to present an in-depth analysis of the ionised gas in the host galaxy of HE~1029$-$1401, but we also perform a tentative analysis of the stellar continuum. Throughout the paper we assume a cosmological model with $H_0=70$\\,km/s, $\\Omega_\\mathrm{m}=0.3$, and $\\Omega_\\Lambda=0.7$. This corresponds to a physical scale of $1.6$\\,kpc/\\arcsec\\ at $z=0.086$. ",
        "conclusions": "We presented optical integral field spectroscopy observations of the luminous radio-quiet QSO HE~1029$-$1401 with VIMOS instrument at the VLT.  A dedicated decomposition technique for IFU data was used to decompose the spectra of the host and the QSO component. This allowed us to perform a detailed spatially-resolved analysis of the ionised gas properties based on emission-line diagnostics and a tentative analysis of the integrated stellar continuum at a S/N level of $\\sim$10. We draw the following conclusions from our present study of this QSO:  \\begin{itemize} \\item We found tentative evidence from the stellar continuum for a blue featureless continuum or a young stellar population ($<$100\\,Myr) embedded in the old stellar population (10\\,Gyr) of the massive elliptical host galaxy. The stellar velocity dispersion is estimated to be $\\sigma_*=320$\\,km/s, although the uncertainty of $\\pm90$\\,km/s is considerable. \\\\  \\item The  black hole mass is estimated from QSO spectrum to be $\\log(M_\\mathrm{BH}) = 8.7\\pm0.3$ accreting at an Eddington ratio of $\\log(L_\\mathrm{Bol}/L_\\mathrm{Edd})=-0.9\\pm0.2$.  Our $M_\\mathrm{BH}$ and $\\sigma_*$ measurements are consistent with the $M_{\\mathrm{BH}}-\\sigma_*$ relation of \\citet{Tremaine:2002} within the uncertainties.\\\\  \\item We detected highly ionised gas in the host galaxy up to a projected distance of at least $\\sim16$\\,kpc that is solely ionised by the QSO radiation. The integrated \\Ox\\ luminosity of this EELR is with $\\log(L(\\Ox)/\\mathrm{[erg/s]}) = 41.78\\pm0.1$ one of the most luminous EELR detect around a radio-quiet QSO. The total ionised gas mass was estimated to $3_{-2}^{+6}\\times10^6 M_{\\sun}$, which should be correct at an order-of-magnitude level only.\\\\  \\item The morphology of the ionised gas appears to be biconical with additional faint knots and arcs at larger distances. The rather narrow bicones may imply more collimated ionisation cones of the AGN than predicted by the unification model for such a luminous QSO. Assuming a disc-like distribution of gas, we prefer to interpret the biconical emission-line image as the projection of the partially illuminated gas disc.\\\\  \\item Comparison of photoionisation models with the observed emission-line ratios of the gas indicated that the metallicity of the gas is significantly lower than expected from the mass-metallicity relation of galaxies. We interpret this as an external origin of the gas and argue that this is most likely due to minor merger(s).\\\\  \\item We found that the ionised gas is rotationally-supported with an additional dispersion-dominated component in the circumnuclear region up to $\\sim$3\\,kpc. The ordered velocity field is clearly distorted on larger scales that either support the scenario of minor mergers of the gas or might be speculatively interpreted as filaments on the surface of an AGN-driven outflowing shell/bubble. However, we did not find clear and conclusive evidence for an AGN outflow.\\\\  \\end{itemize}  Our observations draw a scenario of HE~1029$-$1401 in which the massive elliptical host has recently accreted fresh gas from its environment via minor mergers. This process leaves the host morphology almost unchanged, but the black hole accretion can be `re-juvenated' and a disc of gas is formed. If the accretion event is accompanied by a short burst of star formation, the previous red elliptical galaxy would become bluer moving the galaxy back to the `green valley' in the colour-magnitude diagram. The importance of minor mergers in the late evolution of the early-type galaxies population has recently been highlighted \\citep{Naab:2009,Johansson:2009,Bezanson:2009}. It might also be an important process to explain a luminous AGN phase in early-type galaxies as indicated by our single-object study."
    },
    "1005/1005.2990_arXiv.txt": {
        "abstract": "The compatibility of higher-order Laguerre--Gauss (LG) modes with interferometric technologies commonly used in gravitational wave detectors is investigated. In this paper we present the first experimental results concerning the performance of the LG$_{33}$ mode in optical resonators. We show that the Pound-Drever-Hall error signal for a LG$_{33}$ mode in a linear optical resonator is identical to that of the more commonly used LG$_{00}$ mode, and demonstrate the feedback control of the resonator with a LG$_{33}$ mode. We succeeded to increase the mode purity of a LG$_{33}$ mode generated using a spatial-light modulator from 51\\% to 99\\% upon transmission through a linear optical resonator. We further report the experimental verification that a triangular optical resonator does not transmit helical LG modes. ",
        "introduction": "\\label{sec:intro} The sensitivity of second generation gravitational wave detectors such as Advanced LIGO~\\cite{Harry10} and Advanced Virgo~\\cite{VIR-027A-09} will be limited by the thermal noises of the test masses~\\cite{Rowan05}, therefore the gravitational wave community is involved in research into methods for reducing the effects of this noise. One proposed method for thermal noise reduction is to use so called `flat beams' in the main interferometer, in place of the currently standard fundamental LG$_{00}$ beam~\\cite{Vinet07}. This method is currently investigated for possible upgrades of second generation detectors~\\cite{LSCWP09v4}, as well as for third generation detectors~\\cite{et_punturo2010}. Flat beams provide wider intensity distributions for the same optical losses and can therefore average better over the mirror surface distortions caused by the thermal motions. A number of different flat beam shapes have been proposed, such as mesa beams~\\cite{DAmbrosio03}, conical beams~\\cite{Bondarescu08} and higher-order LG beams~\\cite{Mours06}. One advantage of higher-order LG beams over the other flat beams is their potential compatibility with the currently used spherical mirror surfaces. For this reason we find it worthwhile to further investigate the compatibility of higher-order LG beams with gravitational wave interferometer technology. Some of us recently showed that the potential detection rate of binary neutron star inspiral systems with the Advanced Virgo detector could be increased by a factor of 2.1 if the LG$_{33}$ beam was used in place of the LG$_{00}$ beam~\\cite{Chelkowski09}. In addition to the thermal noise benefits, the wider intensity distributions of higher-order LG beams have been shown to reduce the magnitude of thermal aberrations of optics within the interferometers~\\cite{vinet09}. This would reduce the extent to which thermal compensation systems are relied upon in future detectors to reach design sensitivity. The improved thermal noise characteristics of higher-order LG beams over the fundamental beam also makes the technology a strong candidate for improving the sensitivity of optical clock experiments, which currently also approach the mirror thermal noise limit~\\cite{webster08}.  An investigation using numerical simulations, into the sensing and control signals for length and alignment with a LG$_{33}$ beam in advanced detectors also yielded positive results, indicating that the LG$_{33}$ beam performed as well if not better than the \\tem beam in all of the examined criteria~\\cite{Chelkowski09}. We present the results of an experimental follow-on study into the interferometric performance of higher-order LG beams, in order to assess how much of the potential for sensitivity improvement is realizable in practice. The first crucial test for the interferometric performance of \\lag modes is their compatibility with \\mc technology. We show that the mode cleaner effect works equivalently for the \\lag mode as for the \\tem mode in a linear optical resonator as depicted in figure~\\ref{fig:LGmode-lin-cav}. We also demonstrate the non-compatibility of helical LG modes with three mirror mode cleaners. ",
        "conclusions": "\\label{sec:conclusion} Research into the interferometric performance of LG modes is important for the gravitational wave community, as LG modes offer a thermal noise advantage over the fundamental \\tem mode that can improve the sensitivity achievable by future detectors~\\cite{Mours06}. Simulations have already shown promising results for the interferometric performance of LG modes~\\cite{Chelkowski09}. This article provides experimental verification for some of these predictions: We have demonstrated the generation of a PDH error signal from a linear mode cleaner injected with both helical and sinusoidal \\lag modes equivalent to the error signal obtained with a \\tem mode. We used this error signal to successfully demonstrate longitudinal control of the linear mode cleaner cavity at resonance for the \\lag mode; a vital technique for the operation of gravitational wave interferometers with \\lag modes. We also showed an increase in the purity of a sinusoidal \\lag mode from 51\\% to 99\\% upon transmission through a linear mode cleaner, demonstrating that very high-purity \\lag mode light sources can be produced in this way. Furthermore, we have demonstrated the decomposition of a helical \\lag mode into the constituent sinusoidal \\lag modes with a triangular mode cleaner; a result which has a strong impact on the choice of the optical design of future detectors.  The prospects for LG modes in gravitational wave detectors remain intact following the investigation described in this paper. In the future we will expand this work to use LG modes in systems that combine Michelson interferometers and resonant cavities, and also include alignment control systems. The requirements for mirror surfaces for high finesse systems with LG modes also needs further investigation to move scrutiny of the interferometric performance of LG modes to the next level, towards a possible implementation in future gravitational wave detectors."
    },
    "1005/1005.2397_arXiv.txt": {
        "abstract": "We perform direct numerical simulations of the tidal encounter of a rotating planet on  a highly eccentric or parabolic orbit  about a central star formulated as an initial value problem. This approach enables us to  extend previous work of  Ivanov \\& Papaloizou to consider planet models with solid cores and to avoid making an anelastic approximation.  We obtain a power spectrum of the tidal response of coreless models which enables  global inertial modes to be identified. Their frequencies are found to be in good agreement with those obtained using either a WKBJ approach or the anelastic spectral approach adopted in previous work for small planet rotation rates. We also find that the dependence of the normal mode frequencies on the planet angular velocity in case of higher rotation rates can for the most part be understood   by  applying  first order perturbation theory to the anelastic  modes.  We calculate  the energy and angular momentum exchanged as a result of the tidal encounter and  for coreless models again find good agreement with results obtained using either  the anelastic spectral method.  Models with a solid core  showed evidence of {the emission of shear layers at critical latitudes and possibly } wave attractors after the encounter but the total energy exchanged during the encounter did not differ dramatically from the coreless case  { as long as the ratio of the  core radius to the total radius was less than $50\\%$,  there being hardly any difference at all when  this ratio   was  less than $25\\%$ of the total radius.}  We give a physical and mathematical interpretation of this result.  Finally we are able to validate the use of the anelastic  approximation for both the work presented here and our previous work which led to estimates of  circularisation rates for planets in highly eccentric orbits. ",
        "introduction": "A significant number of extrasolar planets are found in close circular orbits around their central star while others are found in highly eccentric orbits. This is an indication that planet-planet scattering may have been important after formation and generated highly eccentric orbits that subsequently undergo  circularisation  with attendant tidal heating ( see Papaloizou \\& Terquem 2006 for a review and references therein).  In order to investigate such a scenario the tidal interaction of a   planet on a highly eccentric orbit with the central star has to be evaluated. This was considered in the non rotating case by Ivanov \\& Papaloizou (2004) and in the rotating coreless case using an anelastic approximation by Papaloizou \\& Ivanov (2005) and Ivanov \\& Papaloizou (2007) (hereafter referred to as PI and IP, respectively). The excitation of  pulsations or normal modes of the planet is a key aspect of the tidal  problem.   In this paper we extend our previous work by relaxing the anelastic approximation and considering planet models with solid cores. We investigate pulsations of  a uniformly rotating fully convective or barotropic object by  performing numerical simulations of tidal encounters considered  as  initial value problems. The formalism we adopt  does not require  an anelastic approximation  to be made and so it enables  both the response of the low frequency inertial  modes of oscillation as well as the higher frequency $f$ and $p$ modes to be taken into account. As in our preceding studies  we  consider the pulsations to be  excited in an object referred to  hereafter  as a 'planet'  as a consequence of  moving on a parabolic orbit around a central object.  We calculate the energy and angular momentum exchanged after pericentre passage for a variety of pericentre distances and  angular velocities of rotation of the planet. We also study the spectrum of normal modes that is excited and compare their properties with the analytic theory developed in our accompanying paper Ivanov \\& Papaloizou (2009) hereafter  IPN and also results obtained using the anelastic spectral approach described in Ivanov \\& Papaloizou (2007) (hereafter referred to as IP) . This comparison finds good agreement between the different approaches,  validating the use of the anelastic approximation for  describing the inertial mode response.  We also consider the effects of introducing a solid core into the planet. Although we find evidence for the development of singular phenomena such as { the emission of shear layers at critical latitudes and the formation of } wave attractors (see Ogilvie \\& Lin 2004, Ogilvie 2009 { and Rieutord \\& Valdettaro 2010} and references therein) and  after the encounter, the energy  exchanged is  in general not changed by  a very large amount  and  for a modest sized core  hardly at all. We give both a physical and mathematical explanation for this result in appendix B.  The plan of the paper is as follows. In section 2 we review some definitions and basic equations relegating a description of the numerical scheme we use to  solve initial value problems to appendix A. Note that no anelastic approximation is made here. In section 3 we go on to describe the results we obtained for the response of a simple model of a giant planet as a polytrope of index $n=1$ to tidal encounters.  We present results for the total energy and angular momentum transferred to a planet without a solid core for  the full range of rotation rates and  a variety of encounter pericentre distances.  We then go on to consider planet models with solid cores of radius $25\\%$ and $50\\%$ of the total radius illustrating  evidence of the focusing of inertial waves,  {  the emission of shear layers at critical latitudes}  and wave attractors, in these cases. The role of these phenomena has  recently been emphasised by Ogilvie \\& Lin (2004) and Ogilvie (2009) for the situation where the planet is tidally forced while  in a fixed circular orbit. However, in our case with the smaller core, the energy and angular momentum transferred hardy differs from  what is found for the coreless case. In the larger core  case the interaction is somewhat stronger with the energy transfer being about one order of magnitude larger for the  model with a core.  In section 4 we obtain the normal mode frequencies for the coreless planet model by  taking Fourier transforms of the time series  provided by the post tidal interaction responses and locating peaks in the corresponding power spectra. The normal mode frequencies are compared with those obtained using the spectral method described in PI and also the WKBJ method described in IPN, both of which  adopted an anelastic approximation.  The spectra obtained from the simulations are seen to contain  $p$ and  $f$ modes in addition to the inertial modes. The  frequency locations of the main global inertial modes are  found to be in good agreement with those obtained by other methods for azimuthal mode numbers $m=0$ and $m=2.$  In section  4.1.1 we investigate the dependence of the simulation normal mode frequencies (or eigenfrequencies) on the planet rotation rate paying a special attention to the case of inertial waves. In the anelastic approximation it can be shown that eigenfrequencies of inertial modes are proportional to the planet's rotation rate, $\\Omega $, so it is natural to express their values in terms of it. This proportionality breaks down when the  anelastic approximation is not used and the positions of the mode  eigenfrequencies measured in units of $\\Omega $  are  shifted  when it  changes. We compare positions of the simulation normal mode frequencies with those obtained using the anelastic spectral method but with a correction obtained by applying the first order perturbation theory given in IPN. The agreement is very good apart for one case that is apparently affected by an avoided  crossing.   For completeness we consider the tidal response of a coreless model of the type considered by Goodman \\& Lackner(2009) that is in hydrostatic equilibrium under a fixed quadratic gravitational potential in section 5. These authors pointed out that  a simple analytic solution is available and shows that no inertial mode excitation is expected in this case. Our numerical results indicate that  any residual  inertial mode excitation occurring as a result of numerical effects  is very small in this case as should be expected.  In section 6 we compare the total energy and angular momentum transferred to a coreless polytropic planet with the anelastic  spectral results already given in IP. We conclude that these are in good agreement and therefore the anelastic approximation is a valid approximation. Thus we do not confirm the suggestion of Goodman \\& Lackner (2009) that these exchanges could be seriously overestimated as a consequence of its use. Finally in section 7 we discuss our results. ",
        "conclusions": "In this paper we extended  our previous work (IP) on the tidal interaction of a planet on  a highly eccentric or parabolic orbit  about a central star  by relaxing the anelastic approximation and considering  the tidal response of planet models with solid cores. These changes were accomplished by carrying out  numerical simulations of tidal encounters considered  as  initial value problems. The response of  $f$ and $p$ modes  was found in addition to that due to   low frequency inertial  modes.  We calculated  the energy and angular momentum exchanged after pericentre passage for a variety of pericentre distances and  angular velocities of rotation of the planet.  Models both  with and without a solid core were simulated. In the latter case the state variables showed evidence of { the emission of shear layers and/or} wave attractors after the encounter.  In the coreless case we studied  the spectrum of excited inertial  modes and compared it with the spectra obtained in IP and with the analytic theory described in  IPN.  We found good agreement for slow rotation. We were also able to account for the variation of the normal mode frequencies with the planet's angular velocity. In the case of inertial waves a non-trivial variation is induced by terms neglected in the anelastic approximation and it is found to be  small  for rotation frequencies significantly less than critical.   Thus  our analysis indicated that this approximation is generally  a valid one for the inertial modes.  We presented results for the total energy and angular momentum transferred to a planet without a solid core for  the full range of rotation rates and  a variety of encounter pericentre distances and found good agreement in the coreless case with results obtained by IP which were obtained using an anelastic approximation. When a solid core is introduced the energy  exchanged  between star and planet is  in general not changed by  a very large amount. We  give  both a physical and mathematical explanation of this result in appendix  B.  It is  a consequence of the good agreement of our results with those obtained previously using other methods  that the anelastic approximation is valid for the global inertial modes in the slow rotation limit. Thus we do not see the large enhancement of the tidal interaction that Goodman \\& Lackner (2009) indicated would occur as a consequence of adoption of the anelastic approximation.  Since in this limit  our numerical results fully agree with the previous ones obtained in  Ivanov $\\&$ Papaloizou 2004 and IP, where estimates of the circularisation time scales of a planet on a highly eccentric orbit have been made, we can state that these estimates are now confirmed by direct numerical simulations that do not use the anelastic approximation."
    },
    "1005/1005.0392_arXiv.txt": {
        "abstract": "The EAGLE instrument for the E-ELT is a multi-IFU spectrograph, that uses a MOAO system for wavefront correction of interesting lines of sight.  We present a Monte-Carlo AO simulation package that has been used to model the performace of EAGLE, and provide results, including comparisons with an analytical code.  These results include an investigation of the performance of compressed reconstructor representations that have the potential to significantly reduce the complexity of a real-time control system when implemented. ",
        "introduction": "The next generation of optical ground-based \\elts is currently in the design phase, with plans for primary mirror diameters of over 30~m \\cite{tmt,eelt}.  Once built, these facilities will allow astronomers to probe the universe with unprecedented sensitivity and very high resolution.  A suite of instruments for these telescopes is planned, allowing many different observation goals to be met.  The EAGLE instrument for the planned 42~m \\eelt is currently in the design phase \\cite{eagleScience}.  It is a multi-object integral field unit spectrograph using \\ao with a \\moao system to correct incoming wavefronts in open-loop, using wavefront sensors which do not sense the corrections made to the science fields.  Baseline designs for EAGLE include up to 11 wavefront sensors, using laser and natural guide stars.  It is envisaged that there will be 20 science field pick-offs, allowing good \\ao correction in 20 separate fields, each 1.5~arcseconds diameter, simultaneously across a five arc-minute field.  The use of a multi-object \\ao system allows good atmospheric correction to be achieved for selected objects across a wide field of view.  Part of the design phase for EAGLE includes extensive simulation and modelling of the \\ao performance, since \\ao is an essential part of the instrument design.  The simulation and modelling are carried out in two phases.  First, an analytical code is used to obtain an order-of-magnitude performance estimate, covering a large parameter space relatively quickly.  However, many fine details which are essential to include for \\elt scale designs are not included.  A Monte-Carlo code is then used to fill in details giving a more reliable performance estimate, including non-linear effects, and noise sources.  However Monte-Carlo simulation has far greater computational requirements, so a reduced parameter space is considered.  The \\dasp is a Monte-Carlo code which can be used for the simulation of any common form of \\ao system (including classical \\ao, laser tomographic \\ao, multi-conjugate \\ao and \\moao \\cite{2007SPIE.6584E...6D}), and has been developed specifically with \\elt simulation in mind \\cite{basden5}.  It is an end-to-end time-domain code and is parallelised, allowing it to be used across a computing cluster using the \\mpi library to reduce computation time. It includes detailed models of telescope and \\ao systems, allowing high fidelity models to be produced.  The development of a real-time control system for EAGLE is a challenge.  There are expected to be of order $10^5$ wavefront slope measurements, and of order $10^4$ \\dm actuators to control per science path (of which there will be about 20).  These wavefront slope measurements will be used to update the \\dm actuators at about 250~Hz. It is likely that EAGLE will use a conventional matrix-vector based wavefront reconstruction, though other techniques, such as iterative algorithms have not been ruled out.  In this paper, we concentrate only on the matrix-vector based wavefront reconstruction and consider some details that may make this easier to implement in hardware.  In order to access all the elements of the control matrix for each new set of slope measurements, a data rate of order 20~TBs$^{-1}$ is therefore required (assuming four bytes per matrix element).  This will require a very advanced control system, and so any achievable simplifications are desirable.  Here, we present some recent results obtained from the simulation of EAGLE using the \\dasp.  Some of these results are compared with those from an analytical code, produced independently by another member of the EAGLE consortium, where appropriate.  However, the analytical code is unable to include non-linear effects, and so can only be used for rough performance estimates.  The simulations presented here include an investigation of compressed wavefront reconstructor algorithms which could simplify real-time control system design.  The technical difficulty of the real-time control system design is such that it should be considered even at the early design phases.  We also discuss the issue of Shack-Hartmann sensor non-linearities for open-loop systems.  In \\S2, we describe the simulations that have been carried out, in \\S3 we give results, and conclusions are made in \\S4. ",
        "conclusions": "We have performed full end-to-end Monte-Carlo simulations of an \\ao system for EAGLE.  Investigations reported here show that the atmospheric turbulence profile must be well known, with the heights of turbulent layers known to within a few hundred metres.  We have also reported on an investigation of compressed reconstructor representations and find that it is possible to reduce control matrix memory requirements by almost a factor of four in the cases investigated, significantly reducing the complexity of an \\fpga based real-time control system.  An investigation into the effect that the non-linearity of Shack-Hartmann based wavefront sensors has on \\ao system performance has also been carried out, demonstrating that a linearity calibration should be included in an open-loop real-time control system to improve performance."
    },
    "1005/1005.5324_arXiv.txt": {
        "abstract": "We present new {\\it Chandra} X-ray data of the NGC 1333 embedded cluster and combine these data with existing  {\\it Chandra} data, {\\it Sptizer} photometry and ground based spectroscopy of both the NGC 1333 and Serpens Cloud Core clusters to perform a detailed study of the X-ray properties of two of the nearest embedded clusters to the Sun.  We first present new, deeper observations of NGC 1333 with {\\it Chandra ACIS-I}  and combine these with existing {\\it Spitzer} observations of the region. In NGC 1333, a total of 95 cluster members are detected in X-rays, of which 54 were previously identified in the {\\it Spitzer} data.  Of the {\\it Spitzer} identified sources, we detected 23\\% of the Class I protostars, 53\\% of the Flat Spectrum sources, 52\\% of the Class II, and 50\\% of the Transition Disk young stellar objects (YSO).   Forty-one Class III members of the cluster are identified, bringing the total identified YSO population to 178.  The X-ray Luminosity Functions (XLFs) of the NGC 1333 and Serpens clusters are compared to each other and the Orion Nebula Cluster.  Based on a comparison of the XLFs of the Serpens and NGC 1333 clusters to the previously published ONC, we obtain a new distance for the Serpens cluster of $360^{+22}_{-13}$~pc.  Using our previously published spectral types, effective temperatures and bolometric luminosities, we analyze the dependence of the X-ray emission on the  measured stellar properties.   The X-ray luminosity was found to depend on the calculated bolometric luminosity as in previous studies of other clusters.  We examine the dependence of $L_{X}$ on stellar surface area and effective temperature, and find that  $L_{X}$ depends primarily on the stellar surface area.   In the NGC 1333 cluster,  the Class III sources have a somewhat higher X-ray luminosity for a given surface area. We also find evidence in NGC 1333 for a jump in the X-ray luminosity between spectral types of M0 and K7, we speculate that this may result from the presence of radiative zones in the K-stars.  The gas column density vs. extinction in the NGC 1333 parental molecular cloud was examined using the Hydrogen column density determined from the X-ray absorption to the embedded stars and the $K$-band extinction measured to those stars.  In NGC 1333, we find $N_H$ $= 0.89\\pm0.13\\times10^{22}$~$A_K$, this is lower than expected of the standard ISM but similar to that found previously in the Serpens Cloud Core. ",
        "introduction": "In recent years many studies of young stellar clusters have been undertaken to investigate the emission properties of pre-main sequence stars and protostars in the higher energy X-ray regime \\citep{wol,get,prefei}. These studies compliment those carried out in the mid-IR, where young stars are identified by the excess emission produced through the reprocessing of stellar radiation by circumstellar material. Optical and near-IR spectroscopic observations of the young stars in these regions add information on the fundamental properties of the stars, such as effective temperature and luminosity.  Young stellar objects (YSOs) often possess  levels of X-ray emission elevated to $ L_{X} \\sim 10^4 \\times L_{X\\odot}$, this elevated emission can be used to distinguish  them from field stars  \\citep{fei2,fei3}. In developed, hydrogen burning stars, X-ray activity arises from magnetic fields generated as a result of shear between the core radiative zone and the outer convective zone.  The process behind the generation of the highly increased levels of emission in young stellar objects remains uncertain since low mass pre-main sequence stars are often fully convective.  Some suggested causes are magnetic disk-locking between the star and disk \\citep{hay,iso,rom}, accretion onto the star \\citep{kas,fav1,fav2}, and alternative dynamo models for coronal emission \\citep{kuk,giam}.  We present a study of X-ray emitting young stars in the Serpens and NGC 1333 clusters.  These two clusters are examples of nearby ($<$0.5~kpc), low mass regions with high fractions of protostars yet different spatial geometries and frequencies of jets.   Both are deeply embedded in their natal clouds. Serpens is the more deeply embedded with extinctions exceeding 40 magnitudes in the visual.   IR and submm observations identify at least 38 protostars in the central core region, and more than 130 embedded young stars in total \\citep{tes1,tes2,dav1,hog,winston}. An age of $\\le$2~Myrs has been found for the cluster \\citep{kaa,win08a}, with some evidence for a halo of older sources surrounding the central core \\citep[see][for a discussion of the possible age spread.]{win08a} The Chandra X-ray observations of Serpens have previously been reported by \\citet{gia} and \\citet{winston}. NGC 1333 contains a high number of embedded protostars, cores and outflows \\citep{knee}, and while showing a spread in isochronal ages was found to have a median age of 2~Myrs \\citep{win08a}. The region has previously been observed in X-rays with ROSAT by \\citet{prei97} who identify 16 YSOs. A study by \\citet{get2} with {\\it Chandra} associated 95 detections with cluster members and found no difference in X-ray luminosity between Classical T-Tauri stars (CTTS) and weak line T-Tauri stars (WTTS).   \\citet{pre03} observed NGC 1333 using XMM, detecting 86 sources in the region.   Optical and IR surveys to identify the YSOs in the region have been carried out by \\citet{aspin},  \\citet{wilk2},  \\citet{lada},  \\citet{gut07} and \\citet{gut09}. Over 160 young stars have been identified in the cluster.   NGC 1333 does not appear to be centrally condensed though the protostars are observed to trace the underlying gas distribution, whereas the Serpens protostars are located along a filament in the centre of the cluster.  There are four main results in this paper.  First, we utilise the elevated X-ray luminosity to identify YSOs in the NGC 1333 region that do not exhibit IR-emission from a dusty disk (evolutionary Class III) and would otherwise be indistinguishable from field stars using {\\it Spitzer} photometry. Protostars (Class 0/I and Flat Spectrum) and pre-main sequence stars with disks (Class II and Transition Disks) may also show elevated X-ray emission and we identify such objects with detectable X-ray emission.  Second, we use the X-ray data to redetermine the distance to the Serpens cluster.  One of the more elusive properties of Young Stellar Clusters (YSCs) is an accurate measure of their distances. The distance underpins luminosity calculations, and therefore the determination of the stellar isochronal ages and masses.   The distance to NGC 1333 has recently been measured based on a VERA measurement of maser parallax toward SVS 13 to be at 240~pc, and this is the distance used in this study  \\citep{hirota}. A method for determining the distance to YSCs based on the comparison of the X-ray luminosity function (XLF) to a 'universal' XLF \\citep{fei} is applied to the Serpens data to obtain a more accurate distance to the cluster. The most commonly adopted distance to the Serpens Cloud is 260~pc  based on an extinction versus distance diagram to $\\sim$450 stars using Vilnius seven-color photometry \\citep[a discussion of the distance is given in][]{strai}. Other estimates have ranged from 700~pc \\citep{zha} to 440~pc \\citep{racine,strom} and 310~pc \\citep{del}. Here we will attempt to judge the distance via a technique applied to the YSOs in the cluster themselves.  In the third section, we present a detailed study of the combined spectral and X-ray properties of the populations of YSOs as a function of their evolutionary class in the Serpens and NGC 1333 clusters, identified by observations with {\\it Spitzer} and {\\it Chandra} and published by \\citet{gut08} and \\citet{winston,win08a}.  {\\it Chandra} has recently examined the X-ray properties of the Orion Nebula Cluster (ONC), in a project known as the {\\it Chandra} Orion Ultra-Deep Project (COUP).   Although the Serpens and NGC 1333 samples are much smaller than the COUP survey, the analysis of these two smaller clusters has the advantage of excellent 3.6-24~${\\mu}m$ photometry; observations longward of 4.5~${\\mu}m$ are difficult in the centre of the ONC due to the bright infrared emission from the molecular cloud. This photometry allows us to accurately ascertain the evolutionary class of each object, and study the dependence of the X-ray properties on evolutionary state. Further, optical and infrared spectroscopy has been obtained and used to determine the effective temperatures, bolometric luminosities, and stellar surface areas of a number of the YSOs. With these data, we can examine the X-ray flux to bolometric luminosity by evolutionary class, and examine the dependance on surface area and effective temperature in these two clusters.  Finally, we derive the relationship between Hydrogen column density derived from X-ray data and near-infrared extinction. Consistent with our preliminary work, we find that in these molecular clouds, the hydrogen column density per magnitude of K-band extinction is $\\sim$1/3  that in the diffuse ISM.  This may be the result of grain growth and coagulation in cold, dense molecular clouds. ",
        "conclusions": "In this paper we have presented {\\it Chandra} X-ray observations of the NGC 1333 star forming region, and compared these data with the {\\it Spitzer} catalogue of identified young stellar objects in the region. The work further presents a comparative study of the X-ray and fundamental stellar (bolometric luminosity, effective temperature, age, mass) properties of the NGC 1333 and Serpens clusters determined with far-red and near-IR spectra.  \\begin{itemize}  \\item  In total, we identify 95 NGC 1333 cluster members in X-rays. Five of these detections were Class I, eight Flat Spectrum, thirty-nine Class II, and two Transition Disk objects.  We identify 41 Class III members of the NGC 1333 cluster through their elevated X-ray emission, 19 of which are new to this study.  This brings the total number of detected cluster members to 178.  \\item  We found the fraction of each class detected in X-rays; 23\\% of Class I, 53\\% of Flat Spectrum, 52\\% of Class II, 50\\% of Transition Disks.   If the detection rate of Class II and Class III sources are the same, the total Class III population of the cluster is approximately  84 members (implying 221 members in total); however, if the Class III sources are more X-ray luminous than the Class II sources then the number of Class III sources will be lower.  None of the known Class 0 protostars were detected by {\\it Chandra}.  \\item  Using the XLF of Serpens, a new distance to the cluster is estimated at $360^{+22}_{-13}$~pc. This distance is similar to previous literature distances to the cluster, and is found to be compatible with the HRD of the region, giving a median age of 1~Myr for the star forming region.  This new distance may also be applied to the recently discovered Serpens South cluster.  \\item  The $L_X \\propto L_{bol}$ relation found in other star forming regions was confirmed for the two clusters.  This was found to be predominantly a dependence of X-ray luminosity on the stellar surface area, with some weak dependence on the effective temperature.   There is some evidence for a jump in $L_X$ with $T_{eff}$ near the M0-K7 transition, which may be due to the turnover from purely convective YSOs to one with a developing radiative core.  \\item  A marginal dependence of X-ray luminosity with evolutionary class suggests that the diskless Class III sources in NGC 1333 are, for a given stellar surface area,  more luminous than their disk-bearing Class II counterparts. In Serpens the two classes show similar luminosities, although this may be due to the smaller sample size and the larger scatter.  \\item  The temperature of the X-ray emitting plasma was also examined as a function of effective temperature and surface area, with increasing trends observed with surface area.   The Class III sources are found to have lower values of $kT$ than the Class II in each case, with K-S probabilities of them arising from the same distribution of 21\\% (Serpens) and 3.3\\% (NGC 1333).  This may be a result of underestimation of the absorption due to circumstellar disks.  \\item  The variation of X-ray luminosity with age was considered, indicating a marginal trend in NGC 1333 to higher levels of X-ray flux at younger ages.  This may reflect the dependence of isochronal age on stellar surface area. There is also a  trend in $F_X$ with mass, with more massive YSOs showing higher X-ray fluxes.   This is consistent with previous studies such as the {\\it COUP} in Orion.  \\item  In the case of NGC 1333, the ratio of $N_H$ to $A_K$ is lower than expected,  $N_H$ $= 0.89\\pm0.13\\times10^{22}$~$A_K$, a similar value to that found in Serpens \\citep{winston}.   The median value  for both clusters is $N_H$ $= 0.79\\pm0.19\\times10^{22}$~$A_K$. This may be the result of grain growth through coagulation of the accretion of grain mantles as discussed in \\citet{winston}.  \\end{itemize}   This work is based on observations made with the {\\it Chandra} Telescope, under NASA contract NAS8-03060. This work is based on observations made with the {\\it Spitzer} Space Telescope (PID 6, PID 174), which is operated by the Jet Propulsion Laboratory, California Institute of Technology under NASA contract 1407. Support for this work was provided by NASA through contract 1256790 issued by JPL/Caltech. Support for the IRAC instrument was provided by NASA through contract 960541 issued by JPL. This work is based on observations taken with the {\\it Hectospec} instrument on the MMT, a joint venture of the Smithsonian Institute and the University of Arizona. 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 Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation. 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.            \\clearpage"
    },
    "1005/1005.3114_arXiv.txt": {
        "abstract": "We show that a massive black hole binary might resonantly trap a small third body ({\\it e.g.} a neutron star) down to a stage when the binary becomes relativistic due to its orbital decay by gravitational radiation.  The final fate of the third body would be quite interesting for relativistic  astrophysics.  For example,  the parent binary could expel the third body with a velocity more than 10 \\% of the speed of light.  We also discuss the implications of this three-body system  for direct gravitational wave observation. ",
        "introduction": "The orbital period of Pluto is 3/2 times that of Neptune, and their mutual stability is maintained by  this simple commensurability (termed {\\it the mean motion resonance}).  In addition to this well-known fact, the Solar system has various forms of the  mean motion resonances \\cite{ssd}.  Furthermore, more than 8 exoplanet systems are known to have two planets in mean motion resonances.  One of  the notable properties here is that once two planets are trapped in a stable resonance relation, they often keep the state for a long time.  For example,  recent numerical studies showed that some resonant trappings are strong enough to be preserved against planet migration during which orbital radii of two planets  decreased by more than 1 order of magnitude \\cite{Lee:2008nv}.  Black hole (BH) binaries are considered as fascinating sources of broad astrophysical phenomena, and, at the same time, their inspirals and mergers are promising targets for gravitational wave (GW) observation that would also provide us with ideal opportunities to test fundamental aspects of gravity. As in the case of planetary systems, a massive BH binary  might resonantly trap a small body, {\\it e.g.} through interaction with its surrounding disk, and  shrink its orbit by dissipative processes. In this paper we call this kind of three-body system  a resonant multiple inspiral (RMI), and study the evolution and the final fate of an RMI, including the post-Newtonian effects. We find that a BH binary has potential to keep a third body down to a stage when the binary becomes relativistic due to its orbital decay by gravitational radiation.  We also mention impacts of RMIs on relativistic astrophysics and future GW observation. ",
        "conclusions": "Our study for RMIs is based on the 1PN restricted three-body analysis for a circular parent binary.  This simple treatment can be extended in diverse ways. One of the principle directions is to include higher-order PN effects and spins of BHs. A complementary approach would be to inject a test particle into a numerically evolved BH binary (for recent breakthrough see \\cite{Pretorius:2005gq}). From Newtonian analyses we expect that modest eccentricity $(e\\lsim 0.1)$ of the parent binary would not largely change our basic results \\cite{el}, but this should be checked specifically.    Meanwhile, we have regarded the third body as a massless pointlike particle. Actually,  the assumption for the mass  would  be largely relaxed for RMIs \\cite{gascheau}. For example, Saturn has stable co-orbital satellites (in a 1:1 resonance);  Janus and Epimetheus whose masses are comparable (1:0.25) \\cite{ssd}.  Furthermore, the internal structure of the third body might cause interesting astrophysical effects ({\\it e.g.} for  a main-sequence star or even a gas cloud). Beyond a three-body problem,   a parent BH binary might keep many small objects around its $L_4$ and $L_5$ points, similar to the Trojan asteroids of the Sun-Jupiter system. In the case of a 1:1  resonance, we have a severe upper limit $\\mu<\\mu_1$ for the restricted three-body problem, but a larger mass ratio might be allowed for resonances other than 1:1.   In addition to these refinements for  the evolution of RMIs, follow-up studies in related fields are worth exploring. More detailed analyses for detectability of GW signatures at various stages of RMIs would be fruitful, especially for space interferometers such as LISA.  Finally,   discussions on the formation mechanism of RMIs would be an attractive topic on astrophysics (see {\\it e.g.} \\cite{Fujii:2010rd} and references there in for related studies).   The authors would like to thank  H. Arakida and K. Tomida  for their kind support. This work was supported by the Grant-in-Aid for the Global COE Program \"The Next Generation of Physics, Spun from Universality and Emergence\" from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan."
    },
    "1005/1005.1261_arXiv.txt": {
        "abstract": "{Models of the formation, evolution and photoevaporation of circumstellar disks are an essential ingredient in many theories of the formation of planetary systems. The ratio of disk mass over stellar mass in the circumstellar phase of a disk is for a large part determined by the angular momentum of the original cloud core from which the system was formed. While full 3D or 2D axisymmetric hydrodynamical models of accretion onto the disk automatically treat all aspects of angular momentum, this is not so trivial for 1D and semi-2D viscous disk models.} % {Since 1D and semi-2D disk models are still very useful for long-term evolutionary modelling of disks with relatively little numerical effort, we investigate how the 2D nature of accretion affects the formation and evolution of the disk in such models. A proper treatment of this problem requires a correction for the sub-Keplerian velocity at which accretion takes place.} % {We develop an update of our semi-2D time-dependent disk evolution model to properly treat the effects of sub-Keplerian accretion. The new model also accounts for the effects of the vertical extent of the disk on the accretion streamlines from the envelope.} % {The disks produced with the new method are smaller than those obtained previously, but their mass is mostly unchanged. The new disks are a few degrees warmer in the outer parts, so they contain less solid CO\\@. Otherwise, the results for ices are unaffected. The 2D treatment of the accretion results in material accreting at larger radii, so a smaller fraction comes close enough to the star for amorphous silicates to be thermally annealed into crystalline form. The lower crystalline abundances thus predicted correspond more closely to observed abundances than did earlier model predictions. We argue that thermal annealing followed by radial mixing must be responsible for at least part of the observed crystalline material.} % {} % ",
        "introduction": "\\label{sec:intro} With the enormous increase in the amount of high-quality observational data of circumstellar disks in the last few years, a clear picture is emerging of how these objects evolve in time \\citep{jorgensen07a,jorgensen09a,looney07a,lommen08a,sicilia09a}. They form during the collapse of a pre-stellar cloud core, undergo a number of accretion events (FU Orionis and EX Lupi outbursts), live for 3 to 10 Myr, and, shortly before they are destroyed, open up huge gaps visible in the dust continuum and sometimes also in gas lines \\citep{dalessio05a,goto06a,brittain07a,ratzka07a,brown08a,pontoppidan08b}. These physical changes are echoed in the evolution of their chemical composition and dust properties. Pre-stellar cores contain mostly simple hydrides, radicals and other small molecules, largely frozen out onto the cold dust grains \\citep{bergin97b,lee04a}. A fully formed disk is predicted to contain a much richer chemical mixture with a wide variety of complex organic molecules \\citep{rodgers03a,aikawa08a}, although only simple organics have been observed so far \\citep{lahuis06a,carr08a,salyk08a}. The dust by this time has grown from less than a micron to millimetres and centimetres, and part of it has evolved from an amorphous to a crystalline structure \\citep{bouwman01a,bouwman08a,vanboekel05a,natta07a,lommen07a,lommen09a,watson09a,olofsson09a}. Crystalline silicate dust is observed down to temperatures of 100 K, well below the threshold of 800 K required to convert amorphous silicates into crystalline form. One of the central questions of this paper is how much silicate material comes close enough to the star to be crystallised. We also investigate how the crystalline silicates end up so far outside of the hot inner disk where they appear to be formed.  One way to answer these questions is to construct detailed models of the evolution of circumstellar disks based on our current understanding of the physics of these objects, and then compare to the available observational data. However, it would require extraordinarily heavy computations to run a model that does justice to all physical processes known to be involved. A circumstellar disk ranges from a few stellar radii to hundreds of AU, and lasts for several million years. A full model would therefore have to resolve hundreds of millions of inner orbits, and span some five orders of magnitude on a spatial scale. Moreover, an accurate radiative transfer method is required to properly compute the temperatures. All this is clearly too demanding. Most multidimensional hydrodynamical simulations therefore solve sub-problems that only capture part of the disk, or only evolve over a limited time. Even these models, though, require days or weeks of CPU time for a single set of parameters.  An alternative approach is to parameterise most of the physics in some form, and treat the disk evolution as a simpler one-dimensional (1D) time-dependent problem. One assumes axisymmetry and integrates the density vertically to obtain the surface density $\\Sigma$, which is now only a function of the radial coordinate $R$ and the time $t$. These kinds of models go back to the pioneering work by \\citet{shakura73a} and \\citet{lyndenbell74a}. In order to use these models throughout the disk's lifetime, some way must be found to also include the birth phase of the disk in a reasonably realistic way. Two-dimensional axisymmetric hydrodynamical models of disk formation show the presence of a stand-off shock that decelerates the supersonically infalling matter as it approaches the disk from above and below \\citep[e.g.,][]{tscharnuter87a,yorke93a,neufeld94a}. The structure of this stand-off shock is clearly multidimensional in the outer regions, but may be approximated in a simpler manner in the inner regions \\citep{nakamoto94a}. \\citet{hueso05a} constructed a 1D disk evolution model with such an approximation and used it to analyse two T Tauri stars. This showed that simple models of disk formation and evolution can be very powerful and yield valuable insight into the evolutionary stage of young stellar objects. Similar models have also been used to analyse the statistics of the accretion rates measured in pre--main-sequence stars \\citep{dullemond06c,vorobyov08a}.  Another problem addressed with these 1D parameterised models is the origin, evolution and transport of gas and dust in circumstellar disks. For instance, \\citet[hereafter DAW06]{dullemond06a} suggested that the initial outward expansion of the disk during the disk formation phase (observationally the Class 0/I phase) may be very effective in transporting thermally processed dust to the outer parts of the disk. Based on that work, one would expect to find a number of disks with nearly 100\\% crystalline dust. However, no such extremely crystalline disks are observed \\citep{bouwman08a,watson09a}. This is one of the issues addressed in this paper.  Yet another application of parameterised disk evolution models was shown in \\citet[hereafter V09]{visser09a}. We followed the envelope material from thousands of AU inwards, through the accretion shock and into the disk, to analyse when ices evaporate from and recondense onto the grains. Since much of the interesting physics happens in the outer regions of the disk (several hundred AU), where the accretion shock no longer has a simple 1D shape, some way had to be found to include the 2D axisymmetric nature of this region without having to resort to a full-scale multidimensional hydrodynamical simulation. Our recipe was to construct a semi-2D disk model, i.e., to generate the 2D density structure $\\rho(R,z,t)$ (where $z$ is the vertical coordinate away from the midplane) out of the 1D surface density $\\Sigma(R,t)$ using the disk temperature to compute the scale height at every radius. The accretion onto the disk was then modelled by following infalling matter on ballistic supersonic trajectories until its density equals the density of the disk. This yields a 2D axisymmetric shape of the stand-off shock that is similar to that obtained by full 2D axisymmetric hydrodynamical models \\citep{yorke93a,brinch08a}. However, it is not trivial to link this type of multidimensional infall structure back to the 1D disk evolution model. The main obstacle is that a substantial fraction of the matter does not really fall onto the disk's {\\em surface} but onto the disk's {\\em outer edge}, i.e., it accretes from the side instead of from the top. Moreover, this matter rotates with a very sub-Keplerian velocity. If this is not treated in some way, the total angular momentum balance of the system is violated, potentially leading to very wrong estimates of the evolution of the disk mass with time. We devised an ad-hoc solution to this problem in \\vddd{}, which gave reasonable results but did not properly conserve angular momentum. We derive a more rigorous solution in the current paper.  The problem of sub-Keplerian accretion onto a disk is not studied here for the first time. It has long been known that material falling in an elliptic orbit onto the surface of the disk has sub-Keplerian angular momentum, even in the case of infinitely flat disks. When mixing with the disk material, which is in Keplerian rotation, a torque is exerted that pushes the disk material in towards the star. \\citet{cassen81a} described this problem and solved it elegantly. For disks without internal torque and with small vertical extent, they found an analytical solution for the disk surface density as a function of time, and they also presented numerical results for viscous disks. \\citet{hueso05a} derived an alternative solution: they simply calculated the radius for which the specific angular momentum of the infalling matter is Keplerian, and inserted the matter into the disk at that radius. This is not unreasonable, because one may assume that the material, after it hits the surface of the disk, first adjusts its own orbit before it mixes with the disk material below. It may be argued that as long as the angular momentum budget is not violated, the disk evolution is not much affected by the precise treatment. However, if one wishes to follow the radial motion and mixing of certain chemical species or types of dust, this may depend more sensitively on the way the angular momentum problem is treated.  In this paper we add a consistent treatment of sub-Keplerian 2D accretion to a semi-2D model for the formation and viscous evolution of a disk (Sect.\\ \\ref{sec:eqs}). Some basic disk properties arising from the new treatment are discussed in Sect.\\ \\ref{sec:diskprop}, and the results from \\vddd{} are briefly revisited in Sect.\\ \\ref{sec:gasice}. The dust crystallisation results from \\daw{} are re-evaluated in Sect.\\ \\ref{sec:cryst}, and finally conclusions are drawn in Sect.\\ \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc} This paper presents a new method of correcting for the sub-Keplerian velocity of envelope material accreting onto an axi\\-symmetric two-dimensional circumstellar disk. Unlike the previous corrections of \\citet{hueso05a} and \\citet{visser09a}, this new method properly conserves angular momentum and produces infall trajectories without discontinuities. The latter is important for tracing changes in the chemical contents and dust properties during the evolution of the envelope and disk.  The disks produced with the new method are smaller than those produced with the old method by up to a factor of ten. Depending on the initial conditions, the disk masses are between 100 and 200\\% of previously computed values (Sect.\\ \\ref{sec:diskprop}). The new disks are a few degrees colder in the inner regions and a few degrees warmer in the outer regions, resulting in lower abundances of CO ice (Sect.\\ \\ref{sec:gasice}). By the time the system reaches the classical T Tauri stage, at about 1 Myr, the global ice abundances still agree well with observations. Overall, there are no major changes in the gas-ice ratios compared with \\citet{visser09a}.  The disk was treated as geometrically flat by \\citet{dullemond06a}. As in \\citet{visser09a}, we now also take into account the vertical structure when computing the infall trajectories. This results in the bulk of the accretion occuring at larger radii. A smaller fraction of the infalling material now comes close enough to the star to be heated above 800 K, the temperature required for thermal annealing of amorphous silicates into crystalline form. Therefore, the new method produces crystalline abundances that are lower by a few per cent to more than a factor of five compared to the old model. We now obtain a better match with observations and we argue that thermal annealing followed by radial mixing is responsible for at least part of the crystalline silicates observed in disks. An expanded model, which tracks specific forms of crystalline silicate, is required to establish in more detail the importance of this and other possible crystallisation mechanisms."
    },
    "1005/1005.3052_arXiv.txt": {
        "abstract": "We present a {\\em Chandra} monitoring campaign of the highly variable Seyfert galaxy UGC~4203 (the ``Phoenix Galaxy'') which revealed variations in the X-ray absorbing column density on time scales of two weeks. This is the third, clear case, after NGC~1365 and NGC~7582, of dramatic N$_H$ variability on short time scales observed in a ``changing~look'' source, i.e. an AGN observed in the past in both a reflection-dominated and a Compton-thin state. The inferred limits on the distance of the X-ray absorber from the center suggest that the X-ray ``torus'' could be one and the same with the broad emission line region. This scenario, first proposed for an ``ad-hoc'' picture for NGC~1365, may be the common structure of the circumnuclear medium in AGN. ",
        "introduction": "The structure, size and composition of the circumnuclear medium of Active Galactic Nuclei (AGN) is the subject of several studies at different wavelengths. In particular, X-ray absorption variability has proven to be a powerful tool in order to establish the physical properties of the X-ray absorber/reflector: the almost ubiquitous absorption variability detected in bright absorbed AGN (Risaliti et al.~2002) is in itself an indication of a clumpy structure of the absorber. The time scale of such variations provide constraints on the dimensions and the distance from the central black hole of the obscuring clouds, under the assumption of Keplerian motion. The best characterization so far of the physical parameters of the X-ray absorber has been possible for the AGN in NGC~1365. This source has been monitored several times in the last few years with {\\em Chandra}, {\\em XMM-Newton}, {\\em Suzaku}, and showed extreme variability in the X-rays: changes from Compton-thick to Compton-thin states on time scales from weeks (Risaliti et al.~2005) to days (Risaliti et al.~2007) to $\\sim10$~hours (Risaliti et al.~2009A) and eclipses due to Compton-thin clouds ($N_H\\sim$a few $10^{23}$~cm$^{-2}$) on similar time scales (Risaliti et al.~2009B). These observations strongly support the view of an X-ray absorber made of clouds with $N_H$ from $10^{23}$ to $>$10$^{24}$~cm$^{-2}$, densities 10$^{10}$-10$^{11}$~cm$^{-3}$ and distance from the center of hundreds of gravitational radii. These are all values typical of the broad emission line clouds (BELC). The global picture emerging for NGC~1365 is therefore quite different from the standard view of the AGN unified models (Antonucci~1993): here, the X-ray absorber seems not to be associated to the parsec-scale dusty ``torus'' (Krolik \\& Begelman~1988) obscuring the BELC, but to the BELC themselves.  This scenario appears quite well justified for the single case of NGC~1365, whose extreme variability is so far unique among AGN. Two explanations of this uniqueness are possible: (1) NGC~1365 is intrinsically different from (most of) the other AGN; (2) NGC~1365 is an ``average'' AGN with favourable observational conditions (e.g. a particular line of sight through a torus made of clouds of different $N_H$, as in Risaliti et al.~2005).  In order to investigate this issue, the obvious strategy is to look for similar  variations in other obscured AGN.  The most extreme known cases of spectral variations in AGN are those of the ``changing~look'' AGN (Matt et al.~2003), i.e. sources with past X-ray observations in both a Compton-thin and a reflection-dominated state. Members of this class, besides NGC~1365, are NGC~6300 (Guainazzi.~2002), UGC~4203 (Guainazzi et al.~2002), NGC~2992 (Gilli et al.~2000), NGC~7674 (Bianchi et al.~2005), NGC~7582 (Piconcelli et al.~2007). In the latter case, a recent {\\em Suzaku} campaign revealed N$_H$ variations in time scales of a day (Bianchi et al.~2009), making NGC~7582 a second case of a \u2018\u2018changing look'' AGN with fast N$_H$ variability.  The main open question about these sources is whether the observed dramatic spectral changes are due to variability of the intrinsic emission or of the circumnuclear absorber. The case of NGC~2992 is clear: multiple observations on a $\\sim20$~years period strongly suggest an intrinsic fading of the X-ray source (Gilli et al.~2000). We note that this is the only case among the ``changing look'' sources with a column density in the Compton-thin state of the order of 10$^{22}$~cm$^{-2}$, while the other sources have $N_H$$>$$10^{23}$~cm$^{-2}$.  Here we present the results of a {\\em Chandra} monitoring campaign of one of the three remaining sources, UGC~4203 (other names: MKN~1210 and ``Phoenix Galaxy''). Another {\\em Chandra} campaign on one of the last two sources, NGC~6300, will be discussed in a future paper. ",
        "conclusions": ""
    },
    "1005/1005.1100.txt": {
        "abstract": " ",
        "introduction": "Galaxy clusters are the largest gravitationally bound objects in the universe. They are also among the rarest objects in the universe. While these two facts about galaxy clusters may seem disparate, they are in fact intimately related. Our current theory of the origin and evolution of galaxy clusters places them within the broader context of cosmological structure formation in which galaxies, galaxy groups, galaxy clusters, and galaxy superclusters all arise from gravitational instability amplifying perturbations in the cold dark matter density field in an expanding universe. At early times the perturbations are linear in amplitude, and are extremely well described by a Gaussian random field with a known power spectrum (the $\\Lambda CDM$ power spectrum; cf. Fig. \\ref{fig.ps}). At later times, density perturbations become nonlinear and collapse into gravitationally bound systems. The shape of the $\\Lambda CDM$ power spectrum is such that structure forms from the ``bottom up\", with galaxies forming first and clusters forming later. It just so happens that we live in a universe in which cluster-scale perturbations collapsed rather recently (since $z \\sim 1$), which accounts for their rarity as well as their sometimes complex substructure.  As cluster-scale perturbations collapse, they bring in all matter within a sphere of comoving radius of about 15 Mpc, which includes galaxies, intergalactic gas, and anything else in that patch of the universe. Because the escape velocity of galaxy clusters is of order 1000 km/s, everything but relativistic particles become trapped in the cluster potential well. For this reason it is often said that clusters represent a fair sample of the universe. This is true from the standpoint of their matter content. However, from the standpoint of cosmological structure this could not be further from the truth. Galaxy clusters form and evolve in the rarest peaks ($\\sim 3\\sigma$) of the density field (Fig. \\ref{fig.cube}). Galaxy formation begins sooner in such regions, and the galaxies evolve due to internal and external processes which are somewhat different from the general field (e.g., ram pressure stripping).  \\begin{figure} \\begin{center} %\\includegraphics[width=\\textwidth]{Norman-f1.eps} \\includegraphics[width=\\textwidth]{cube.eps} \\end{center} \\caption{AMR hydrodynamic cosmological simulation of cosmic structure in a 700 Mpc volume of the universe. Up to seven levels of adaptive mesh refinement (AMR) resolve the distribution of baryons within and between galaxy clusters, for an effective resolution of $65,536^3$. Volume rendering of baryon density. From \\cite{Norman07}.} \\label{fig.cube} \\end{figure}   Galaxy clusters are interesting objects in their own right, and for decades have been extensively studied in the optical, Xray, and radio wavebands \\cite{Carnegie04}. More recently this has been extended to the microwave, infrared, and extreme UV \\cite{Taiwan03,Varenna04}, motivated in part by the fact that galaxy clusters are excellent cosmological probes. Because of their large size and high X-ray luminosities, they can be seen to great distances. As discussed in this volume and in \\cite{Varenna04}, galaxy clusters can also be seen in absorption/emission against the cosmic microwave background (CMB) via the Sunyaev-Zeldovich effect (SZE). As discussed by Rephaeli elsewhere in these proceedings, the SZE is redshift independent, meaning that deep microwave surveys should detect all galaxy clusters in a particular region of the sky regardless of their redshift provided the telescope has enough angular resolution and sensitivity. Indeed such surveys are underway now and results are expected soon.  One of the most challenging measurements in modern observational cosmology is the dark energy equation of state which describes the time rate of change of the vacuum energy density $\\dot{\\Lambda}$ responsible for the accelerating expansion of the universe \\cite{Perlmutter03}. Galaxy clusters were identified by the US Dark Energy Task Force (DETF) as one of four complementary methods for doing this. However in order to measure $\\dot{\\Lambda}$ we must measure the cluster mass distribution function versus redshift to very high accuracy. Accurately measuring the mass of a galaxy cluster is actually quite difficult despite the number of ingenious techniques that have been developed. This motivates attempts to measure the ``mass-observable\" relationships in observed samples of clusters, and find the ones with the least scatter and least bias.  Numerical simulations are helpful in this regard, as one can in principle calibrate the mass-observable relationships by comparing simulated observations with {\\em in situ} measurements. Where this has been done, a discouraging result is found: most of the observables of a given simulated cluster depend sensitively on numerical resolution and assumed baryonic physics, including radiative cooling, star formation, galactic and AGN feedback processes. However it has been shown that the integrated SZE signal is rather insensitive to assumed baryonic physics \\cite{Motl04,Nagai06}, perhaps relaxing requirements on modelers and giving encouragement to SZE cluster observers that the dark energy program may be feasible after all. The point is that the astrophysics of galaxy clusters and their utility as cosmological probes are inextricably linked, and both are worthwhile of study.  I was asked by the organizers to lecture on three topics relavant to the theme of the summer school. First I was asked to review the standard cosmological framework and basic results from the theory of cosmological structure formation within which galaxy clusters can be understood. Second, I was asked to review how galaxy clusters are simulated on a computer, and summarize the basic findings. I do this in the next two sections, which are slightly updated and abbreviated versions of the lecture notes I published in the 2004 Varenna Summer School volume \\cite{Varenna04}. Finally I was asked to give a lecture of my choosing, which was on recent progress in galaxy cluster modeling focusing on the incorporation of additional baryonic physics and simulating SZE surveys . These topics are presented in Sections 5 and 6 respectively.  In line with the character of the summer school, I have attempted to be pedagogical, emphasizing the key concepts and results that a student needs to know if s/he wants to understand the current literature or do research in this area. Literature citations are kept to a minimum, except for textbooks, reviews, and research papers that I found to be particularly helpful in preparing this article. I am indebted to Dr. Rocky Kolb whose slides much of Section 2 are based upon.  \\section {Cosmological framework and perturbation growth in the linear regime}  Our modern theory of the structure and evolution of the universe, along with the observational data which support it, is admirably presented in the textbook by Dodelson \\cite{dodelson03}. Remarkable observational progress has been made in the past two decades which has strengthened our confidence in the correctness of the hot, relativistic, expanding universe model (Big Bang), has measured the universe's present mass-energy contents and kinematics, and lent strong support to the notion of a very early, inflationary phase. Moreover, observations of high redshift supernovae unexpectedly have revealed that the cosmic expansion is accelerating at the present time, implying the existence of a pervasive, dark energy field with negative pressure \\cite{Perlmutter03}. This surprising discovery has enlivened observational efforts to accurately measure the cosmological parameters over as large a fraction of the age of the universe as possible, especially over the redshift interval 0 $<$ z $<$ 1.5 which, according to current estimates, spans the deceleration-acceleration transition. These efforts include large surveys of galaxy large scale structure, galaxy clusters, weak lensing, the Lyman alpha forest, and high redshift supernovae, all of which span the relevant redshift range. Except for the supernovae, all other techniques rely on measurements of cosmological structure in order to deduce cosmological parameters.   \\subsection{Cosmological standard model} The dynamics of the expanding universe is described by the two Friedmann equations derived from Einstein's theory of general relativity under the assumption of homogeneity and isotropy. The expansion rate at time $t$ is given by \\begin{equation}\\label{eq1} H^2(t)\\equiv \\left( {\\frac{\\dot {a}}{a}} \\right)^2=\\frac{8\\pi G}{3}\\sum\\limits_i {\\rho _i } -\\frac{k}{a^2}+\\frac{\\Lambda }{3} \\end{equation} where $H(t)$ is the Hubble parameter and $a(t)$ is the FRW scale factor at time $t$. The first term on the RHS is proportional to the sum over all energy densities in the universe $\\rho _{i }$ including baryons, photons, neutrinos, dark matter and dark energy. We have explicitly pulled the dark energy term out of the sum and placed it in the third term assuming it is a constant (the cosmological constant). The second term is the curvature term, where $k=0,\\pm 1$ for zero, positive, negative curvature, respectively. Equation (\\ref{eq1}) can be cast in a form useful for numerical integration if we introduce $\\Omega $ parameters: \\begin{equation}\\label{eq2} \\Omega _i \\equiv \\frac{8\\pi G}{3H^2}\\rho _i ,\\mbox{ }\\Omega _\\Lambda \\equiv \\frac{8\\pi G}{3H^2}\\rho _\\Lambda =\\frac{\\Lambda }{3H^2},\\mbox{ }\\Omega_k \\equiv \\frac{-k}{(aH)^2} \\end{equation} Dividing equation (\\ref{eq1}) by $H^2$ we get the sum rule 1=$\\Omega _{m}+\\Omega _{k}+\\Omega _{\\Lambda }$, which is true at all times, where $\\Omega _{m}$ is the sum over all $\\Omega _{i}$ excluding dark energy. At the present time $H(t)=H_{0}, a=1$, and cosmological density parameters become \\begin{equation}\\label{eq3} \\Omega _i (0)=\\frac{8\\pi G}{3H_0^2 }\\rho _i (0),\\mbox{ }\\Omega _\\Lambda (0)=\\frac{\\Lambda }{3H_0^2 },\\mbox{ }\\Omega _k (0)=\\frac{-k}{H_0^2 } \\end{equation} Equation (\\ref{eq1}) can then be manipulated into the form \\begin{equation}\\label{eq4} \\dot {a}=H_0 [\\Omega _m (0)(a^{-1}-1)+\\Omega _\\gamma (0)(a^{-2}-1)+\\Omega _\\Lambda (0)(a^2-1)+1]^{1/2} \\end{equation} Here we have explicitly introduced a density parameter for the background radiation field $\\Omega _{\\gamma }$ and used the fact that matter and radiation densities scale as a$^{-3}$ and a$^{-4}$, respectively, and we have used the sum rule to eliminate $\\Omega _{k}$. Equation (\\ref{eq4}) is equation (\\ref{eq1}) expressed in terms of the \\textit{current} values of the density and Hubble parameters, and makes explicit the scale factor dependence of the various contributions to the expansion rate. In particular, it is clear that the expansion rate is dominated first by radiation, then by matter, and finally by the cosmological constant.  Current measurements of the cosmological parameters by different techniques \\cite{Komatsu09} yield the following numbers [(0) notation suppressed]: \\[ \\begin{array}{l} h\\equiv H_0 /(100km/s/Mpc)\\approx 0.72 \\\\ \\Omega _{total} \\approx 1,\\mbox{ }\\Omega _\\Lambda \\approx 0.73\\mbox{, }\\Omega _m =\\Omega _{cdm} +\\Omega _b \\approx 0.27,\\Omega _k \\approx 0 \\\\ \\Omega _b \\approx 0.04,\\mbox{ }\\Omega _\\nu \\approx 0.005,\\mbox{ }\\Omega _\\gamma \\approx 0.00005 \\\\ \\end{array} \\] This set of parameters is referred to as the concordance model \\cite{bops99}, and describes a spatially flat, low matter density, high dark energy density universe in which baryons, neutrinos, and photons make a negligible contribution to the large scale dynamics. Most of the matter in the universe is cold dark matter (CDM) whose dynamics is discussed below. As we will also see below, baryons and photons make an important contribution to shaping of the matter power spectrum despite their small contribution to the present-day energy budget. Understanding the evolution of baryons in nonlinear structure formation is essential to interpret X-ray and SZE observations of galaxy clusters.  The second Friedmann equation relates the second time derivative of the scale factor to the cosmic pressure $p $ and energy density\\textit{ $\\rho $} \\begin{equation}\\label{eq5} \\frac{\\ddot {a}}{a}=-\\frac{4\\pi G}{3}(\\rho +3p),\\mbox{ }\\rho =\\sum\\limits_i {\\rho _i } =\\rho _m +\\rho _\\gamma +\\rho _\\Lambda \\end{equation} $p$ and $\\rho $ are related by an equation of state $p_{i}=w_{i}\\rho _{i}$, with $w_{m}$=0, $w_{\\gamma }$=1/3, and $w_{\\Lambda }= -1$. We thus have \\begin{equation}\\label{eq6} \\frac{\\ddot {a}}{a}=-\\frac{4\\pi G}{3}(\\rho _m +2\\rho _\\gamma -2\\rho _\\Lambda )\\mbox{. } \\end{equation} Expressed in terms of the current values for the cosmological parameters we have \\begin{equation}\\label{eq7} \\frac{\\ddot {a}}{a}=-\\frac{1}{2}H_0^2 [\\Omega _m (0)a^{-3}+2\\Omega _\\gamma (0)a^{-4}-2\\Omega _\\Lambda (0)]\\mbox{. } \\end{equation} Evaluating equation \\ref{eq7} using the concordance parameters, we see the universe is currently accelerating $\\ddot {a}\\approx 0.6H_0^2 \\mbox{ }$ . Assuming the dark energy density is a constant, the acceleration began when \\begin{equation}\\label{eq8} a\\equiv \\frac{1}{1+z}=\\left( {\\frac{\\Omega _m (0)}{2\\Omega _\\Lambda (0)}} \\right)^{\\mbox{1/3}}\\mbox{ }\\approx 0.57 \\end{equation} or $z\\sim 0.75$.  \\subsection{The Linear power spectrum}  Cosmic structure results from the amplification of primordial density fluctuations by gravitational instability. The power spectrum of matter density fluctuations has now been measured with considerable accuracy across roughly four decades in scale. Figure \\ref{fig.ps} shows the latest results, taken from reference \\cite{tegmark03}. Combined in this figure are measurements using cosmic microwave background (CMB) anisotropies, galaxy large scale structure, weak lensing of galaxy shapes, and the Lyman alpha forest, in order of decreasing comoving wavelength. In addition, there is a single data point for galaxy clusters, whose current space density measures the amplitude of the power spectrum on 8 h$^{-1}$ Mpc scales \\cite{wef93}. Superimposed on the data is the predicted $\\Lambda $CDM linear power spectrum at z=0 for the concordance model parameters. As one can see, the fit is quite good. In actuality, the concordance model parameters are determined by fitting the data. A rather complex statistical machinery underlies the determination of cosmological parameters, and is discussed in Dodelson (2003, Ch. 11). The fact that modern CMB and LSS data agree over a substantial region of overlap gives us confidence in the correctness of the concordance model. In this section, we define the power spectrum mathematically, and review the basic physics which determines its shape. Readers wishing a more in depth treatment are referred to references \\cite{dodelson03,kolbturner90}.  \\begin{figure}[htbp] \\centerline{\\includegraphics[width=5.18in,height=4.75in]{fig1.eps}} \\caption{Linear matter power spectrum P(k) versus wavenumber extrapolated to z=0, from various measurements of cosmological structure. The best fit $\\Lambda $CDM model is shown as a solid line (from \\cite{tegmark03}.)} \\label{fig.ps} \\end{figure}  At any epoch $t$ (or $a$ or $z)$ express the matter density in the universe in terms of a mean density and a local fluctuation: \\begin{equation}\\label{eq9} \\rho (\\vec {x})=\\bar {\\rho }(1+\\delta (\\vec {x})) \\end{equation} where $\\delta (\\vec {x})$is the density contrast. Expand $\\delta (\\vec {x})$ in Fourier modes: \\begin{equation}\\label{eq10} \\delta (\\vec {x})\\equiv \\frac{\\rho (\\vec {x})-\\bar {\\rho }}{\\bar {\\rho }}=\\int {\\delta (\\vec {k})\\exp (-i\\vec {k}\\cdot \\vec {x})d^3} k. \\end{equation} The autocorrelation function of $\\delta (\\vec {x})$ defines the power spectrum through the relations \\begin{equation}\\label{eq11} \\left\\langle {\\delta (\\vec {x})\\delta (\\vec {x})} \\right\\rangle =\\int\\limits_0^\\infty {\\frac{dk}{k}} \\frac{k^3\\left| {\\delta ^2(\\vec {k})} \\right|}{2\\pi ^2}=\\int\\limits_0^\\infty {\\frac{dk}{k}} \\frac{k^3P(k)}{2\\pi ^2}=\\int\\limits_0^\\infty {\\frac{dk}{k}} \\Delta ^2(k) \\end{equation} where we have the definitions \\begin{equation}\\label{eq12} P(k)\\equiv \\left| {\\delta ^2(\\vec {k})} \\right|,\\mbox{ and }\\Delta ^2(k)\\equiv \\frac{k^3P(k)}{2\\pi ^2}. \\end{equation} The quantity $\\Delta ^2(k)$ is called the dimensionless power spectrum and is an important function in the theory of structure formation. $\\Delta ^2(k)$ measures the contribution of perturbations per unit logarithmic interval at wavenumber $k$ to the variance in the matter density fluctuations. The $\\Lambda $CDM power spectrum asymptotes to $P(k)\\sim k^{1}$ for small $k$, and $P(k)\\sim k^{-3}$ for large $k$, with a peak a $k^{\\star}\\sim 2\\times 10^{-2}$ h Mpc$^{-1}$ corresponding to $\\lambda^{\\star}\\sim $350 h$^{-1}$ Mpc. $\\Delta ^2(k)$ is thus asymptotically flat at high $k$, but drops off as $k^{4}$ at small $k$. We therefore see that most of the variance in the cosmic density field in the universe at the present epoch is on scales $\\lambda <  \\lambda^{\\star}.$  \\begin{figure}[htbp] \\centerline{\\includegraphics[width=4in,height=3in]{fig2.eps}} \\caption{The tale of two fluctuations. A fluctuation which is superhorizon scale at matter-radiation equality grows always, while a fluctuation which enters the horizon during the radiation dominated era stops growing in amplitude until the matter dominated era begins.} \\label{fig2} \\end{figure}  What is the origin of the power spectrum shape? Here we review the basic ideas. Within the inflationary paradigm, it is believed that quantum mechanical (QM) fluctuations in the very early universe were stretched to macroscopic scales by the large expansion factor the universe underwent during inflation. Since QM fluctuations are random, the primordial density perturbations should be well described as a Gaussian random field. Measurements of the Gaussianity of the CMB anisotropies \\cite{komatsu03} have confirmed this. The primordial power spectrum is parameterized as a power law $P_p (k)\\propto k^n$, with $n=1$ corresponding to scale-invariant spectrum proposed by Harrison and Zeldovich on the grounds that any other value would imply a preferred mass scale for fluctuations entering the Hubble horizon. Large angular scale CMB anisotropies measure the primordial power spectrum directly since they are superhorizon scale. Observations with the WMAP satellite yield a value very close to $n=1$ \\cite{Komatsu09}.  To understand the origin of the spectrum, we need to understand how the amplitude of a fluctuation of fixed comoving wavelength $\\lambda$ grows with time. Regardless of its wavelength, the fluctuation will pass through the Hubble horizon as illustrated in Fig. \\ref{fig2}. This is because the Hubble radius grows linearly with time, while the proper wavelength a$\\lambda $ grows more slowly with time. It is easy to show from Eq. \\ref{eq1} that in the radiation-dominated era, $a\\sim t^{1/2}$, and in the matter-dominated era (prior to the onset of cosmic acceleration) $a\\sim t^{2/3}$. Thus, inevitably, a fluctuation will transition from superhorizon to subhorizon scale. We are interested in how the amplitude of the fluctuation evolves during these two phases. Here we merely state the results of perturbation theory (e.g., Dodelson 2003, Ch. 7).    \\begin{figure}[htbp] \\begin{center} \\begin{tabular}{c} \\includegraphics[width=4in,height=2in]{fig3a.eps} \\\\ \\includegraphics[width=4in,height=2in]{fig3.eps} \\end{tabular} \\end{center} \\caption{ a) Evolution of the primordial power spectrum on superhorizon scales during the radiaton dominated era. b) Scale-free spectrum produces a constant contribution to the density variance per logarithmic wavenumber interval entering the Hubble horizon (no preferred scale) c) resulting matter power spectrum, super- and sub-horizon. Figures courtesy Rocky Kolb.} \\label{fig3} \\end{figure}  \\subsection{Growth of fluctuations in the linear regime }  To calculate the growth of superhorizon scale fluctuations requires general relativistic perturbation theory, while subhorizon scale perturbations can be analyzed using a Newtonian Jeans analysis. We are interested in scalar density perturbations, because these couple to the stress tensor of the matter-radiation field. Vector perturbations (e.g., fluid turbulence) are not sourced by the stress-tensor, and decay rapidly due to cosmic expansion. Tensor perturbations are gravity waves, and also do not couple to the stress-tensor. A detailed analysis for the scalar perturbations yields the following results. In the \\underline {radiation dominated era}, \\[ \\begin{array}{l} \\delta _+ (t)=\\delta _+ (t_i )(t/t_i )\\mbox{ superhorizon scales} \\\\ \\delta _+ (t)=constant                \\mbox{ ~~~subhorizon scales} \\\\ \\end{array} \\] while in the \\underline {matter dominated era}, \\[ \\begin{array}{l} \\delta _+ (t)=\\delta _+ (t_i )(t/t_i )^{2/3}\\mbox{ superhorizon scales} \\\\ \\delta _+ (t)=\\delta _+ (t_i )(t/t_i )^{2/3}\\mbox{ subhorizon scales} \\\\ \\end{array} \\] This is summarized in Fig. \\ref{fig2}, where we consider two fluctuations of different comoving wavelengths, which we will call large and small. The large wavelength perturbation remains superhorizon through matter-radiation equality (MRE), and enters the horizon in the matter dominated era. Its amplitude will grow as $t$ in the radiation dominated era, and as $t^{2/3}$ in the matter dominated era. It will continue to grow as $t^{2/3}$ after it becomes subhorizon scale. The small wavelength perturbation becomes subhorizon before MRE. Its amplitude will grow as $t$ while it is superhorizon scale, remain constant while it is subhorizon during the radiation dominated era, and then grow as $t^{2/3}$ during the matter-dominated era.  Armed with these results, we can understand what is meant by a scale-free primordial power spectrum (the Harrison-Zeldovich power spectrum.) We are concerned with perturbation growth in the very early universe during the radiation dominated era. Superhorizon scale perturbation amplitudes grow as $t$, and then cease to grow after they have passed through the Hubble horizon. We can define a Hubble wave number $k_H \\equiv 2\\pi /R_H \\propto t^{-1}.$ Fig. 4a shows the primordial power spectrum at three instants in time for k$<$k$_{H}$. We see that the fluctuation amplitude at k=k$_{H}$(t) depends on primordial power spectrum slope n. The scale-free spectrum is the value of n such that $\\Delta ^2(k_{H}(t))$=constant for k$>$k$_{H}$. A simple analysis shows that this implies n=1. Since $\\Delta ^2(k)\\propto k^3P(k)$, we then have \\[ \\begin{array}{l} P(k)\\propto k^1,\\mbox{ }k\\le k_H \\\\ P(k)\\propto k^{-3},\\mbox{ }k>k_H \\\\ \\end{array} \\]  In actuality, the power spectrum has a smooth maximum, rather than a peak as shown in Fig. 4c. This smoothing is caused by the different rates of growth before and after matter-radiation equality. The transition from radiation to matter-dominated is not instantaneous. Rather, the expansion rate of the universe changes smoothly through equality, as given by Eq. 1, and consequently so do the temporal growth rates. The position of the peak of the power spectrum is sensitive to the time when the universe reached matter-radiation equality, and hence is a probe of $\\Omega _\\gamma /\\Omega _m $.  Once a fluctuation becomes sub-horizon, dissipative processes modify the shape of the power spectrum in a scale-dependent way. Collisionless matter will freely stream out of overdense regions and smooth out the inhomogeneities. The faster the particle, the larger its free streaming length. Particles which are relativistic at MRE, such as light neutrinos, are called hot dark matter (HDM). They have a large free-streaming length, and consequently damp the power spectrum over a large range of k. Weakly Interacting Massive Particles (WIMPs) which are nonrelativistic at MRE, are called cold dark matter (CDM), and modify the power spectrum very little (Fig. \\ref{fig4}). Baryons are tightly coupled to the radiation field by electron scattering prior to recombination. During rcombination, the photon mean-free path becomes large. As photons stream out of dense regions, they drag baryons along, erasing density fluctuations on small scales. This process is called Silk damping, and results in damped oscillations of the baryon-photon fluid once they become subhorizon scale. The magnitude of this effect is sensitive to the ratio of baryons to collisionless matter, as shown in Fig. \\ref{fig4}.  \\begin{figure}[htbp] \\includegraphics[width=2.5in,height=1.7in]{fig4.eps} \\includegraphics[width=2.5in,height=1.7in]{fig5.eps} \\caption{Effect of dissipative processes on the evolved power spectrum. Left: Effect of collisionless damping (free streaming) in the dark matter. Right: Effect of collisional damping (Silk damping) in the matter-radiation fluid. Figures courtesy Rocky Kolb.} \\label{fig4} \\end{figure}  \\section {Analytic models for nonlinear growth, virial scaling \\\\ relations, and halo statistics}  Here we introduce a few concepts and analytic results from the theory of structure formation which underly the use of galaxy clusters as cosmological probes. These provide us with the vocabulary which pervades the literature on analytic and numerical models of galaxy cluster evolution. Material in this section has been derived from three primary sources: Padmanabhan (1993) \\cite{pad93} for the spherical top-hat model for nonlinear collapse, Dodelson (2003) \\cite{dodelson03} for Press-Schechter theory, and Bryan {\\&} Norman (1998) \\cite{BN98} for virial scaling relations.   \\subsection{Nonlinearity defined}  In the linear regime, both super- and sub-horizon scale perturbations grow as $t^{2/3}$ in the matter-dominated era. This means that after recombination, the linear power spectrum retains its shape while its amplitude grows as $t^{4/3}$ before the onset of cosmic acceleration (Fig. 6a). When $\\Delta ^2(k)$ for a given k approaches unity linear theory no longer applies, and some other method must be used to determine the fluctuation's growth. In general, numerical simulations are required to model the nonlinear phase of growth because in the nonlinear regime, the modes do not grow independently. Mode-mode coupling modifies both the shape and amplitude of the power spectrum over the range of wavenumbers that have gone nonlinear.  At any given time, there is a critical wavenumber which we shall call the nonlinear wavenumber k$_{nl}$ which determines which portion of the spectrum has evolved into the nonlinear regime. Modes with k$<$k$_{nl}$ are said to be linear, while those for which k$>$ k$_{nl}$ are nonlinear (Fig. 6b). Conventionally, one defines the nonlinear wavenumber such that $\\Delta (k_{nl} ,z)=1.$ From this one can derive a nonlinear mass scale $M_{nl} (z)=\\frac{4\\pi }{3}\\bar {\\rho }(z)\\left( {\\frac{2\\pi }{k_{nl} }} \\right)^3$. A more useful and rigorous definition of the nonlinear mass scale comes from evaluating the amplitude of mass fluctuations within spheres or radius R at epoch z. The enclosed mass is $M=\\frac{4\\pi }{3}\\bar {\\rho }(z)R^3.$ The mean square mass fluctuations (variance) is \\begin{equation}\\label{eq17} \\left\\langle {(\\delta M/M)^2} \\right\\rangle \\equiv \\sigma ^2(M)=\\int {d^3kW_T^2 (kR)P(k,z),} \\end{equation} where W is the Fourier transform of the top-hat window function \\begin{equation}\\label{eq18} \\begin{array}{l} \\mbox{W(}{\\rm {\\bf x}}\\mbox{)}=\\left\\{ {{\\begin{array}{*{20}c} {3/4\\pi R^3,\\mbox{ }\\left| {\\rm {\\bf x}} \\right|<R} \\hfill \\\\ {0,\\mbox{ }\\left| {\\rm {\\bf x}} \\right|\\ge R} \\hfill \\\\ \\end{array} }} \\right. \\\\ \\to W_T (kR)=3\\left[ {\\sin (kR)/kR-\\cos (kR)} \\right]/(kR)^2. \\\\ \\end{array} \\end{equation} If we approximate P(k) locally with a power-law $P(k,z)=D^2(z)k^m$, where D is the linear growth factor, then $\\sigma ^2(M)\\propto D^2R^{-(3+m)}\\propto D^2M^{-(3+m)/3}.$ From this we see that the RMS fluctuations are a decreasing function of M. At very small mass scales, m$\\rightarrow -3$, and the fluctuations asymptote to a constant value. We now define the nonlinear mass scale by setting $\\sigma $(M$_{nl})$=1. We get that (\\cite{white94}) \\begin{equation}\\label{eq19} M_{nl} (z)\\propto D(z)^{6/(3+m)}\\mbox{ (}\\propto \\mbox{(1}+\\mbox{z)}^{\\mbox{-6/(3}+\\mbox{m)}}\\mbox{ for EdS).} \\end{equation} For $m > -3$, the smallest mass scales become nonlinear first. This is the origin of hierarchical (``bottom-up'') structure formation (Fig. 6b).  \\begin{figure} \\begin{center} \\includegraphics[width=2.5in]{lingrowth.eps} \\includegraphics[width=2.5in]{nonlinearscale.eps} \\end{center} \\caption{Two ways of looking at the growth of post-recombination matter fluctuations in the linear regime. Left: 3D matter power spectrum increases uniformly proportional to the linear growth factor $D(z)$. Measurements are from SDSS galaxy large scale structure data. Right: the evolution of the dimensionless power spectrum $\\Delta^2 (k)$. Nonlinearity is defined where $\\Delta^2 (k_{nl})=1.$ All scales with $k \\geq k_{nl}$ have collapsed into bound objects, and do so in a ``bottom-up\" fashion.} \\label{fig.growth} \\end{figure}  \\subsection{Spherical Top-Hat Model}  \\begin{figure}[htbp] \\centerline{\\includegraphics[width=3in,height=2in]{fig6.eps}} \\caption{Evolution of a top-hat perturbation in an EdS universe. Depending on the E, the first integral of motion, the fluctuation collapses (E$<$0), continues to expand (E$>$0), or asymptotically reaches it maximum radius (E=0). Virialization occurs when the fluctuation has collapsed to half its turnaround radius.} \\label{fig5} \\end{figure}  We now ask what happens when a spherical volume of mass M and radius R exceeds the nonlinear mass scale. The simplest analytic model of the nonlinear evolution of a discrete perturbation is called the spherical top-hat model. In it, one imagines as spherical perturbation of radius $R$ and some constant overdensity $\\bar {\\delta }=3M/4\\pi \\bar{\\rho} R^3$ in an Einstein-de Sitter (EdS) universe. By Birkhoff's theorem the equation of motion for R is \\begin{equation}\\label{eq20} \\frac{d^2R}{dt^2}=-\\frac{GM}{R^2}=-\\frac{4\\pi G}{3}\\bar {\\rho }(1+\\bar {\\delta })R \\end{equation} whereas the background universe expands according to Eq. \\ref{eq6} \\begin{equation}\\label{eq21} \\frac{d^2a}{dt^2}=-\\frac{4\\pi G}{3}\\bar {\\rho }a. \\end{equation}  Comparing these two equations, we see that the perturbation evolves like a universe of a different mean density, but with the same initial expansion rate. Integrating Eq. \\ref{eq20} once with respect to time gives us the first integral of motion: \\begin{equation}\\label{eq22} \\frac{1}{2}\\left( {\\frac{dR}{dt}} \\right)^2-\\frac{GM}{R}=E, \\end{equation} where E is the total energy of the perturbation. If E$<$0, the perturbation is bound, and obeys \\begin{equation}\\label{eq23} \\frac{R}{R_m}=\\frac{(1-cos \\theta)}{2}, ~~~\\frac{t}{t_m}=\\frac{(\\theta-sin\\theta)}{\\pi} \\end{equation} where $R_m$ and $t_m$ are the radius and time of ``turnaround''. At turnaround (as $\\theta \\rightarrow \\pi$), the fluctuation reaches its maximum proper radius (see Fig. \\ref{fig5}). As $t\\rightarrow 2t_m, R\\rightarrow 0$, and we say the fluctuation has collapsed.  A detailed analysis of the evolution of the top-hat perturbation is given in Padmanabhan (1993, Ch. 8) for general $\\Omega_m$. Here we merely quote results for an EdS universe. The mean \\textit{linear} overdensity at turnaround; i.e., the value one would predict from the linear growth formula $\\delta \\sim t^{2/3}$, is 1.063. The actual overdensity at turnaround using the nonlinear model is 4.6. This illustrates that nonlinear effects set in well before the amplitude of a linear fluctuation reaches unity. As R$\\rightarrow $0, the nonlinear overdensity becomes infinite. However, the linear overdensity at $t=2t_m$ is only 1.686. As the fluctuation collapses, other physical processes (pressure, shocks, violent relation) become important which establish a gravitationally bound object in virial equilibrium before infinite density is reached. Within the framework of the spherical top-hat model, we say virialization has occurred when the kinetic and gravitational energies satisfy virial equilibrium: $\\left| U \\right|=2K.$ It is easy to show from conservation of energy that this occurs when $R=R_m/2$; in other words, when the fluctuation has collapsed to half its turnaround radius. The nonlinear overdensity at virialization $\\Delta _c$ is not infinite since the radius is finite. For an EdS universe, $\\Delta _c =18\\pi ^2\\approx 180$. Fitting formulae for non-EdS models are provided in the next section.  \\subsection{Virial Scaling Relations} The spherical top-hat model can be scaled to perturbations of arbitrary mass. Using virial equilibrium arguments, we can predict various physical properties of the virialized object. The ones that interest us most are those that relate to the observable properties of gas in galaxy clusters, such as temperature, X-ray luminosity, and SZ intensity change. Kaiser \\cite{kaiser86} first derived virial scaling relations for clusters in an EdS universe. Here we generalize the derivation to non-EdS models of interest. In order to compute these scaling laws, we must assume some model for the distribution of matter as a function of radius within the virialized object. A top-hat distribution with a density $\\rho =\\Delta _c \\bar {\\rho }(z)$ is not useful because it is not in mechanical equilibrium. More appropriate is the isothermal, self-gravitating, equilibrium sphere for the collisionless matter, whose density profile is related to the one-dimensional velocity dispersion \\cite{bt87} \\begin{equation}\\label{eq24} \\rho (r)=\\frac{\\sigma ^2}{2\\pi Gr^2}. \\end{equation} If we define the virial radius r$_{vir}$ to be the radius of a spherical volume within which the mean density is $\\Delta _{c}$ times the critical density at that redshift ($M=4\\pi r_{vir}^3 \\rho _{crit} \\Delta _c /3)$, then there is a relation between the virial mass M and $\\sigma $: \\begin{equation} \\label{eq25} \\sigma =M^{1/3}[H^2(z)\\Delta _c G^2/16]^{1/6}\\approx 476f_\\sigma \\left( {\\frac{M}{10^{15}M_\\odot }} \\right)^{1/3}(h^2\\Delta _c E^2)^{1/6}\\mbox{ km s}^{\\mbox{-1}}. \\end{equation} Here we have introduced a normalization factor $f_{\\sigma}$ which will be used to match the normailization from simulations. The redshift dependent Hubble parameter can be written as $H(z)=100hE(z)\\mbox{ km s}^{-1}$ with the function $E^2(z)=\\Omega _m (1+z)^3+\\Omega _k (1+z)^2+\\Omega _\\Lambda $, where the $\\Omega$'s have been previously defined.  The value of $\\Delta_c$ is taken from the spherical top-hat model, and is 18$\\pi ^{2}$ for the critical EdS model, but has a dependence on cosmology through the parameter $\\Omega (z)=\\Omega _m (1+z)^3/E^2(z).$ Bryan and Norman (1998) provided fitting formulae for $\\Delta_c$ for the critical  for both open universe models and flat, lambda-dominated models \\begin{equation}\\label{eq26} \\Delta _c =18\\pi ^2+82x-39x^2\\mbox{ for }\\Omega _k =0,\\mbox{ }\\Delta _c =18\\pi ^2+60x-32x^2\\mbox{ for }\\Omega _\\Lambda =0 \\end{equation} where x=$\\Omega $(z)-1.  If the distribution of the baryonic gas is also isothermal, we can define a ratio of the ``temperature'' of the collisionless material ($T_\\sigma =\\mu m_p \\sigma ^2/k)$ to the gas temperature: \\begin{equation} \\label{eq27} \\beta =\\frac{\\mu m_p \\sigma ^2}{kT} \\end{equation} Given equations (\\ref{eq26}) and (\\ref{eq27}), the relation between temperature and mass is then \\begin{equation} \\label{eq28} kT=\\frac{GM^{2/3}\\mu m_p }{2\\beta }\\left[ {\\frac{H^2(z)\\Delta _c }{2G}} \\right]^{1/3}\\approx 1.39f_T \\left( {\\frac{M}{10^{15}M_\\odot }} \\right)^{2/3}(h^2\\Delta _c E^2)^{1/3}\\mbox{ keV,} \\end{equation} where in the last expression we have added the normalization factor f$_{T}$ and set $\\beta $=1.  The scaling behavior for the object's X-ray luminosity is easily computed by assuming bolometric bremsstrahlung emission and ignoring the temperature dependence of the Gaunt factor: $L_{bol} \\propto \\int {\\rho ^2} T^{1/2}dV\\propto M_b \\rho T^{1/2}.$ where M$_{b}$ is the baryonic mass of the cluster. This is infinite for an isothermal density distribution, since $\\rho $ is singular. Observationally and computationally, it is found that the baryon distribution rolls over to a constant density core at small radius. A procedure is described in Bryan and Norman (1998) which yields a finite luminosity: \\begin{equation} \\label{eq29} L_{bol} =1.3\\times 10^{45}\\left( {\\frac{M}{10^{15}M_\\odot }} \\right)^{4/3}(h^2\\Delta _c E^2)^{7/6}\\mbox{ }\\left( {\\frac{\\Omega _b }{\\Omega _m }} \\right)^2\\mbox{ erg s}^{-1}. \\end{equation} Eliminating M in favor of T in Eq. \\ref{eq29} we get \\begin{equation} \\label{eq30} L_{bol} =6.8\\times 10^{44}\\left( {\\frac{kT/f_T }{1.0\\mbox{ keV}}} \\right)^2(h^2\\Delta _c E^2)^{1/2}\\mbox{ }\\left( {\\frac{\\Omega _b }{\\Omega _m }} \\right)^2\\mbox{ erg s}^{-1}. \\end{equation} The scaling of the SZ ``luminosity'' is likewise easily computed. If we define L$_{SZ}$ as the integrated SZ intensity change: $L_{SZ} =\\int {dA\\int {n_e \\sigma _T } } \\left( {\\frac{kT}{m_e c^2}} \\right)dl\\propto M_b T$, then \\begin{equation}\\label{eq30a} L_{SZ} =\\frac{GM^{5/3}\\sigma _T }{2\\beta m_e c^2}\\left[ {\\frac{H^2(z)\\Delta _c }{2G}} \\right]^{1/3}\\left( {\\frac{\\Omega _b }{\\Omega _m }} \\right). \\end{equation} We note that cosmology enters these relations only with the combination of parameters $h^2\\Delta _c E^2$, which comes from the relation between the cluster's mass and the mean density of the universe at redshift z. The redshift variation comes mostly from E(z), which is equal to (1+z)$^{3/2}$ for an EdS universe.  \\subsection{Statistics of hierarchical clustering: Press-Schechter theory} Now that we have a simple model for the nonlinear evolution of a spherical density fluctuation and its observable properties as a function of its virial mass, we would like to estimate the number of virialized objects of mass M as a function of redshift given the matter power spectrum. This is the key to using surveys of galaxy clusters as cosmological probes. While large scale numerical simulations can and have been used for this purpose (see below), we review a powerful analytic approach by Press and Schechter \\cite{ps74} which turns out to be remarkably close to the numerical results. The basic idea is to imagine smoothing the cosmological density field at any epoch z on a scale R such that the mass scale of virialized objects of interest satisfies $M=\\frac{4\\pi }{3}\\bar {\\rho }(z)R^3.$ Because the density field (both smoothed and unsmoothed) is a Gaussian random field, the probability that the mean overdensity in spheres of radius R exceeds a critical overdensity $\\delta _{c}$ is \\begin{equation}\\label{eq31} p(R,z)=\\frac{2}{\\sqrt {2\\pi } \\sigma (R,z)}\\int\\limits_{\\delta _c }^\\infty {d\\delta } \\exp \\left( {-\\frac{\\delta ^2}{2\\sigma ^2(R,z)}} \\right) \\end{equation} where $\\sigma(R,z)$ is the RMS density variation in spheres of radius R as discussed above. Press and Schechter suggested that this probability be identified with the fraction of particles which are part of a nonlinear lump with mass exceeding M if we take $\\delta _c =1.686,$ the linear overdensity at virialization. This assumption has been tested against numerical simulations and found to be quite good \\cite{wef93} (however, see below). The fraction of the volume collapsed into objects with mass between $M$ and $M+dM$ is given by $(dp/dM)dM$. Multiply this by the average number density of such objects $\\rho _m /M$ to get the number density of collapsed objects between $M$ and $M+dM$: \\begin{equation}\\label{eq32} dn(M,z)=-\\frac{\\bar {\\rho }}{M}\\frac{dp(M(R),z)}{dM}dM. \\end{equation} The minus sign appears here because p is a decreasing function of M. Carrying out the derivative using the fact that $dM/dR=3M/R,$ we get \\begin{equation}\\label{eq33} \\frac{dn(M,z)}{dM}=\\sqrt {\\frac{2}{\\pi }} \\frac{\\bar {\\rho }\\delta _c }{3M^2\\sigma }e^{-\\delta _c^2 /2\\sigma ^2}\\left[ {-\\frac{d\\ln \\sigma }{d\\ln R}} \\right]. \\end{equation} The term is square brackets is related to the logarithmic slope of the power spectrum, which on the mass scale of galaxy clusters is close to unity. Eq. \\ref{eq33} is called the \\textit{halo mass function}, and it has the form of a power law multiplied by an exponential. To make this more explicit, approximate the power spectrum on scales of interest as a power law as we have done above. Substituting the scaling relations for $\\sigma $ in Eq. \\ref{eq33} one gets the result \\cite{white94} \\begin{equation}\\label{eq34} \\frac{dn}{dM}=\\left( {\\frac{2}{\\pi }} \\right)^{1/2}\\frac{\\bar {\\rho }}{M^2}\\left( {1+\\frac{m}{3}} \\right)\\left[ {\\frac{M}{M_{nl} (z)}} \\right]^{\\frac{m-3}{6}}\\exp \\left[ {-\\left( {\\frac{M}{M_{nl} (z)}} \\right)^{\\frac{3+m}{3}}/2} \\right]. \\end{equation} Here, $M_{nl} (z)$ is the nonlinear mass scale. To be more consistent with the spherical top-hat model, it satisfies the relation $\\sigma (M_{nl} ,z)=\\delta _c $; i.e., those fluctuations in the smoothed density field that have reached the linear overdensity for which the spherical top-hat model predicts virialization.  \\begin{figure} \\begin{center} \\includegraphics[width=\\textwidth]{hubblevolume.eps} \\end{center} \\caption{The Hubble Volume Simulation \\cite{HubbleVolume}. Left: slice through the dark matter density field. Yellow/orange peaks correspond to galaxy clusters. Right: Halo multiplicity function, as measured from the simulation (solid lines), with Press-Schechter prediction superposed (dashed line). From \\cite{Jenkins01}.} \\label{fig.hubblevolume} \\end{figure}  \\subsection{Validating the Halo Mass Function using N-body Simulations} Below we discuss how one goes about numerically simulating the nonlinear evolution of the density fluctuations described by the $\\Lambda CDM$ power spectrum. Here we simply mention two works which made detailed comparisons of the PS formula with halo populations found in dark matter N-body simulations. The first is by Jenkins et al. (2001)\\cite{Jenkins01} who analyzed the results of the ``Hubble Volume\" simulation--a simulation of dark matter clustering carried out in a cubic volume 3 Gpc/h on a side with $1024^3$ dark matter particles (Fig. \\ref{fig.hubblevolume}). %The growth in the power of parallel computers over the past 10 years, and %simulation codes that can take advantage of that, opens the possibility %of simulating large statistical samples of galaxy clusters all at once. %This approach was pioneered by the Virgo Consortium's Hubble Volume simulation %\\cite{HubbleVolume} (cf. Fig. CC). They simulated a box 3 Gpc/h on a side with %$1024^3$ dark matter particles, This yields a dark matter particle mass of $2.2 \\times 10^{12} M_{\\odot}$, implying that a galaxy cluster halo would typically contain $100-1000$ particles. The relatively poor mass resolution is offset by the very large volume, which permits exploring the cluster mass function across a broad range of masses including the very high mass end. Fig. \\ref{fig.hubblevolume} shows a slice through the simulation volume on which the dark matter density field is plotted. Jenkins et al. (2001) identified dark matter halos using the friends-of-friends algorithm \\cite{Davis85} and found that while the PS formula gives a good approximation to the numerical data, it underpredicts the number of rare, massive objects, and overpredicts the number of ``typical\" objects as shown in the right panel of Fig. \\ref{fig.hubblevolume}.  Warren et al. (2006)\\cite{Warren06} were interested in testing the validity of the PS formula over a wider range of mass scales than can be obtained from a single simulation. They simulated 16 boxes of different physical size but the same number of DM particles ($1024^3$) nested in such a way that together they define a composite halo mass function covering 5 orders of magnitude in mass scale. They derived a fitting formula for the composite halo mass function for the WMAP3 concordance cosmological parameters by assuming a parameterized form for the {\\em halo multiplicity function} of: \\begin{equation} f(\\sigma) = A(\\sigma^{-a}+b)\\exp^{-c/\\sigma^2} \\end{equation} where the multiplicity function $f(\\sigma)$ is related to the mass function $n(M)$ via \\begin{equation} f(\\sigma) = \\frac{M}{\\bar{\\rho}} \\frac{dn}{d ln \\sigma^{-1}} \\end{equation} where $A, a, b$ and $c$ come from the fitting procedure, and are documented in \\cite{Warren06}.  To illustrate one application of this formula, Fig. \\ref{fig.massfn} shows the cluster halo mass function computed using the Warren et al. (2006) fit for three different redshifts for the cosmological parameters adopted in the simulation shown in Fig. 1. Overplotted on the semi-analytic predictions are the halo mass functions obtained from the simulation itself. The departure of the simulation from the predictions at the low mass end are due to finite resolution effects discussed in Sec. 4 below.  \\begin{figure}[htbp] \\begin{center} \\includegraphics[width=4.0in]{Hallman07f3.eps} \\end{center} \\caption{Cumulative number of galaxy clusters exceeding mass $m$ for three redshifts. Solid lines: $z=0.1$, short-dashed lines $z=1$, long-dashed lines $z=2$. The thick lines are from the simulation shown in Fig. 1, while the thin lines are predictions using the Warren \\cite{Warren06} fitting function. Note the rapid redshift evolution of the number of massive clusters. The departure of the simulation from the predictions at the low mass end are due to finite resolution effects. From \\cite{Hallman07}.} \\label{fig.massfn} \\end{figure}    \\subsection{Application to galaxy clusters}    \\begin{figure}[htbp] \\includegraphics[width=2.5in]{Hallman07f2a.eps} \\includegraphics[width=2.5in]{Hallman07f2b.eps} \\caption{Predicted aerial density of galaxy clusters for the WMAP3 concordance cosmological model. Left: varying the amplitude of the matter power spectrum $\\sigma_8$ holding all other parameters fixed; Right: varying the matter density $\\Omega_m$ holding all other parameters fixed. From \\cite{Hallman07}.} \\label{fig.clusters} \\end{figure}  The aerial density of galaxy clusters can be calculated by multiplying the redshift-dependent halo mass function using the Warren fitting formula with the redshift-dependent differential volume element for one square degree on the sky. Fig. \\ref{fig.clusters}a shows the result varying the amplitude of the matter power spectrum $\\sigma_8$ holding all other cosmological parameters fixed, while Fig. \\ref{fig.clusters}b shows the result varying the matter density $\\Omega_m$ holding all other cosmological parameters fixed. In general we see a rapid rise in the number of clusters with increasing redshift due to the increasing volume element. However, the space density of clusters declines rapidly with redshift (see Fig. \\ref{fig.massfn}), and thus the aerial density peaks at $z \\approx 1$ and then declines rapidly toward higher redshift.  From these curves and the virial scaling relations given above it is easy to predict the expected number of clusters of a given X-ray temperature, X-ray luminosity, or SZ luminosity as a function of redshift \\cite{ecf96,BN98}, bearing in mind that real clusters may not perfectly obey the virial scaling relations. In fact they don't as discussed in my 2004 Varenna lectures \\cite{Norman04} and by Borgani in these proceedings. ",
        "conclusions": ""
    },
    "1005/1005.0784_arXiv.txt": {
        "abstract": "{} {Supernova remnants (SNRs) are believed to be the main sources of Galactic cosmic rays. Molecular clouds associated with SNRs can produce gamma-ray emission through the interaction of accelerated particles with the concentrated gas. The middle aged SNR W28, for its associated system of dense molecular clouds, provides an excellent opportunity  to test this hypothesis.}% {We present the AGILE/GRID observations of SNR W28,  and  compare them with observations at other wavelengths (TeV and $^{12}$CO $(J=1\\rightarrow0)$ molecular line  emission). }% { The gamma-ray flux detected by AGILE from the dominant source associated with W28 is  (14 $\\pm$ 5) $\\times$ 10$^{-8}$ ph cm$^{-2}$ s$^{-1}$ for $E > 400$ MeV. This source is positionally well correlated with the TeV emission observed by the HESS  telescope. The local variations of the GeV to TeV flux ratio suggest a difference between the CR spectra of the north-west and south molecular cloud complexes. % A model based on a hadronic-induced interaction and diffusion with two molecular clouds at different distances from the W28 shell can explain both the morphological and spectral features observed by AGILE in the MeV-GeV energy range and by the HESS telescope in the TeV energy range. {The combined set of AGILE and H.E.S.S. data strongly support a hadronic model for the gamma-ray production in W28.}} {}% ",
        "introduction": "SNR W28, also known as G6.4-0.1, is a middle aged supernova remnant (with an age of at least 35 thousand years), located at  a distance between 1.8 and 3.3 kpc in the inner region of the Galaxy (\\textit{l,b} = 6.71, -0.05). W28 is a mixed-morphology SNR, with large (about 48 arcmin) apparent dimensions (see \\cite{green09}). At radio wavelength a shell structure, where the interaction of the SNR ejecta and the ISM creates a shock, is clearly visible (see Figure \\ref{radio}). The spectral index between 328 and 1415 MHz is $\\alpha=-0.35 \\pm 0.18$ and shows large local variations correlated with the flux density (\\cite{dubner00} ). The $^{12}$CO $(J=1\\rightarrow0)$ molecular line observations taken by the NANTEN telescope reveals a system of molecular clouds  associated with the SNR (\\cite{mizuno04}). Contours in Figure \\ref{radio} show the intensity of the $^{12}$CO line emission integrated over the velocity range 3 to 27 km s$^{-1}$ corresponding to the dynamical distances compatible with W28. The emission appears concentrated in two main complexes near the north-west (cloud N) and south (cloud S) part of the SNR. Evidence of interaction between the remnant and the system of molecular clouds is given  {by the detection of 1720 MHz OH {maser} emission (\\cite{frail94})} and by the observation of an  {unusually} high value of the ratio CO $ {(J=3-2)/(J =2-1)}$ (\\cite{arikawa99}). \\\\Two TeV sources have been detected by the HESS Cherenkov telescope near W28 (\\cite{aharonian06}). The sources HESS J1801-233  and HESS J1800-240 (A, B and C) are very well correlated with, respectively, the N and S molecular clouds  resolved by the NANTEN telescope (\\cite{aharonian08}). In the MeV-GeV energy range a source compatible with W28 was reported in the Third EGRET catalog, 3EG J1800-2338 (\\cite{hartman99}). The flux measured by EGRET above 100 MeV was $61.3 \\pm 6.7$ ph cm$^{-2}$ s$^{-1}$, with a photon spectral index of 2.10. In both the Fermi LAT Bright Source List (0FGL J1801.6-2327, see \\cite{fermi_catalog}) and the First AGILE Catalog (1AGL J1803-2258, see \\cite{agile_catalog}) a gamma-ray source centered on the position of cloud N is reported. \\\\We present in this Letter the results of deep AGILE observation of  1AGL J1803-2258 which, at energies greater than 400 MeV, has a shape that turns out to be incompatible with a single point source, and is remarkably correlated with both the molecular cloud system seen by NANTEN and the TeV sources detected by HESS. \\begin{figure}[ht!] \\begin{center} \\includegraphics[width=9cm]{radio_co5.eps} \\caption{Map of VLA 90 cm radio emission from SNR W28  { in celestial J2000 coordinates R.A. and {Dec}. (Colours indicate the intensity in Jy beam$^{-1}$)}. The black contours show the CO intensity emission obtained with the NANTEN radio telescope which trace molecular clouds, integrated over the velocity interval of 3-27 km s$^{-1}$ (contours vary between 90 and 170 K deg with a step size of 10 K deg). Two concentrations of CO emission are clearly visible near the positions $\\alpha, \\delta$ = 270.4, -23.4 (cloud N) and  $\\alpha, \\delta$ = 270.2, -24.1  (cloud S). } \\end{center} \\label{radio} \\end{figure} \\begin{figure}[ht!] \\begin{center} \\includegraphics[width=9cm]{w28_tev3.eps} \\end{center} \\caption{Sky map   {in celestial J2000 coordinates, R.A. and {Dec}.,} of the significance   {($\\sigma$)} of the TeV  emission detected by the HESS telescope (\\cite{aharonian08}). The blue circle indicates the location of the supernova remnant W28. As in Figure \\ref{radio} the black contours show the CO intensity emission. } \\label{tev} \\end{figure} ",
        "conclusions": "The AGILE deep observations of SNR W28 determine the existence of several contributions to the gamma-ray emission from this SNR. A highly significant gamma-ray source (7 $\\sigma$ for E$>$ 100 MeV) is associated with W28. This source has an average  flux of (40 $\\pm$ 11) $\\times$ 10$^{-8}$  ph cm$^{-2}$ s$^{-1}$ during the period covered by the AGILE observations. AGILE also detected, for photon energies above 400 MeV, a SE extension corresponding to a massive molecular cloud. We proposed a model based on a hadronic-induced interaction with two molecular clouds adjacent to the SNR in order to fit the observations. This model explains the morphological and spectral features detected both by AGILE i n the MeV-GeV energy range and by the HESS telescope in the TeV energy range. Setting the distances and masses of the two main molecular clouds to values compatible with the radio CO observations we can estimate that the total energy in protons is $3.3 \\times 10^{49}$ erg. {Our data and analysis provides strong support to a hadronic origin of the gamma-ray emission from W28.}"
    },
    "1005/1005.2923_arXiv.txt": {
        "abstract": "{ The goal of the {\\it Herschel} Open Time Key programme {\\it ``TNOs are Cool!\"} is to derive the physical and thermal properties for a large sample of Centaurs and trans-Neptunian objects (TNOs), including resonant, classical, detached and scattered disk objects. We present results for seven targets either observed in PACS point-source, or in mini scan-map mode. {\\it Spitzer}-MIPS observations were included for three objects. The sizes of these targets range from 100\\,km to almost 1000\\,km, five have low geometric albedos below 10\\%, \\object{(145480) 2005 TB$_{190}$} has a higher albedo above 15\\%. Classical thermal models driven by an intermediate beaming factor of $\\eta$=1.2 or $\\eta$-values adjusted to the observed colour temperature fit the multi-band observations well in most cases. More sophisticated thermophysical models give very similar diameter and albedo values for thermal inertias in the range 0-25\\,J\\,m$^{-2}$\\,s$^{-0.5}$\\,K$^{-1}$, consistent with very low heat conductivities at temperatures far away from the Sun. The early experience with observing and model strategies will allow us to derive physical and thermal properties for our complete {\\it Herschel} TNO sample of 140 targets as a benchmark for understanding the solar system debris disk, and extra-solar ones as well.} ",
        "introduction": "Trans-Neptunian objects (TNOs) are believed to represent one of the most primordial populations in the solar system (Morbidelli et al.\\ \\cite{morbidelli08}). The TNO population comprises (i) the main Kuiper belt beyond the orbit of Neptune ($\\sim$32 - 50\\,AU), consisting of objects in resonant and non-resonant orbits, and (ii) the halo outskirts of ``scattered\" and ``detached\" bodies beyond 50\\,AU. The Centaurs, an unstable orbital class of minor planets (e.g., Horner et al.\\ \\cite{horner03}; \\cite{horner04}), are closer to the Sun and in transition from the Kuiper belt towards the inner solar system. More than 1300 TNOs have been detected so far, revealing a rich orbital structure and intriguing physical properties. The Trans-Neptunian population is analogous to the debris disks observed around several other, 5-500\\,Myr old stars (Moro-Martin et al.\\ \\cite{moro08}, Jewitt et al.\\ \\cite{jewitt09}). This analogy is bolstered by similarities in sizes and observed masses (typically 30-300\\,AU and 0.01-0.1\\,M$_{\\oplus}$  for the ``exo-disks\"), with the important difference that the detected mass in extra-solar debris disks is in the form of $\\sim$10-1000\\,$\\mu$m, short-lived, dust particles. The vast majority of the mass in these disks is invisible to us, probably in the form of kilometre (or more)-sized bodies, resembling trans-Neptunian objects.  As part of the {\\it Herschel} (Pilbratt et al.\\ \\cite{pilbratt10}) Science Demonstration Phase we observed 17 targets in different instrument configurations and observing modes. Here we present the analysis of early photometric measurements with the Photodetector Array Camera and Spectrometer (PACS - Poglitsch et al.\\ \\cite{poglitsch10}) of five TNOs and two Centaurs. The science aspects from longer wavelengths Spectral and Photometric Imaging Receiver (SPIRE - Griffin et al.\\ \\cite{griffin10}) observations on \\object{(136472) Makemake} and \\object{(90482) Orcus} are included in Lim et al.\\ (\\cite{lim10}) and the thermal lightcurve of \\object{(136108) Haumea} is presented by Lellouch et al.\\ (\\cite{lellouch10}). The full Open Time Key Programme includes about 140 TNOs (M\\\"uller et al.\\ \\cite{mueller09}). ",
        "conclusions": "Our small and dynamically very inhomogeneous sample confirms the consistency between different model techniques and nicely agrees with {\\it Spitzer} results (Stansberry et al.\\ \\cite{stansberry08}) on three overlap targets. Based on the seven targets we also showed the model capabilities for multiple, dual, single or even non-detections. Models with either a beaming parameter of $\\sim$1.2 or thermal inertias below 25\\,J\\,m$^{-2}$\\,s$^{-0.5}$\\,K$^{-1}$ explain the measured SEDs, confirming low heat conductivities at temperatures far away from the Sun. The target sizes range from diameters below 100\\,km to almost 1000\\,km. The derived geometric albedos are below 10\\%. Only \\object{(145480) 2005 TB$_{190}$} has an albedo above 15\\%, which is possibly related to its very unusal orbit located entirely outside the major planets. The derived low albedos are close to the mean value of 8\\% given by Stansberry et al.\\ (\\cite{stansberry08}). Albedo trends with object size, dynamic or taxonomic types are not yet visible in the small sample."
    },
    "1005/1005.2330_arXiv.txt": {
        "abstract": "The supermassive black holes harboured in active galactic nuclei are at the origin of powerful jets which can emit copious amounts of $\\gamma$-rays. The exact interplay between the infalling matter, the black hole and the relativistic outflow is still poorly known, and this parallel session of the $12^{\\rm th}$ Marcel Grossman meeting intended to offer the most up to date status of observational results with the latest generation of ground and space-based instruments, as well as the theoretical developments relevant for the field. ",
        "introduction": "The detection of cosmic $\\gamma$-rays is essentially based on either direct or indirect methods. The former uses satellised high-energy particle physics type detectors, based on the tracking of the $\\gamma$-ray pair conversion and the calorimetry of the subsequent electromagnetic cascade. This technique is constrained almost by definition in weight, size and power consumption, which limits respectively the maximal $\\gamma$-ray energy, the effective area, and the dynamic range. They have however a large field of view (fov), almost avoiding any need for pointed observations since the performance of the instrument (angular and energy resolution) changes slowly within the fov. With livetimes close to 90\\%, space-based observatories provide some of the best continuous monitoring of the high-energy sky ever realized.  The indirect methods rely on the detection of the secondary particles induced by the $\\gamma$-ray pair creation in the atmosphere. This can be done directly on the ground, with large area all-weather water \\v{C}erenkov detectors or air shower detectors (such as MILAGRO and ARGO-YBJ) at the price of $\\sim$TeV energy thresholds and poor background rejection. Another technique consist in the indirect detection of the air cascade at energies $\\geq10\\,{\\rm GeV}$, using the faint optical \\v{C}erenkov flash of the shower with atmospheric \\v{C}erenkov telescopes (ACTs), with the drawback of having a limited livetime $\\sim 10\\%$ due to the same constraint as ground-based optical telescopes of clear, dark night skies for carrying out observations limiting light curves to be continuous only over a time up to $\\sim 6{\\rm h}$, with varying sensitivities and energy thresholds due to the specifics of this technique. Since such cascades, or even the \\v{C}erenkov light pool, have transverse sizes $o(100\\,{\\rm m})$, the effective area can be much larger than the physical size of the detector itself, and, provided an efficient hadronic background, have a tremendous sensitivity. Clearly, space- and ground-based instruments are extremely complementary, and are very likely to remain so until much larger $\\gamma$-ray detectors can be put in space.   \\subsection{Space missions}  After a long generation of $\\gamma$-ray detectors based on moderately precise spark-chambers technology running on expandable gas, such as SAS-2 (1972-1973), COS-B (1975-1982) and EGRET (on board of \\textit{CGRO}, 1991-2000), the newest and most recent observatories \\textit{AGILE} and \\textit{Fermi}-LAT use robust and finely segmented silicon-based trackers, improving the angular resolution by an order of magnitude\\footnote{For a summary of the \\textit{Fermi}-LAT performance see {\\em http://www-glast.slac.stanford.edu/software/IS/glast\\_lat\\_performance.htm}} but also increasing the fov up to about $\\pi\\,{\\rm sr}$ through a more compact design. the LAT onboar \\textit{Fermi} detects $\\gamma$-rays in the range $20\\,{\\rm MeV}$--$300\\,{\\rm GeV}$ where the lowest energy is limited by the background and the highest energy mostly by statistics. The whole sky is covered within two orbits, or $3\\,{\\rm h}$, with an integrated flux sensitivity above $100\\,{\\rm MeV}$ of $\\sim 2\\times 10^{-6}\\,{\\rm cm^{-2}}\\,{\\rm s^{-1}}$ and $\\sim 4\\times 10^{-7}\\,{\\rm cm^{-2}}\\,{\\rm s^{-1}}$ for respectively Galactic and extragalactic ones, providing among the longest light curves ever recorded\\cite{tosti}\\!.  An extremely important figure is that the LAT detects the Crab nebula, a standard candle used in TeV astronomy, within a year at $\\sim100\\,{\\rm GeV}$\\cite{fermicrab}\\!, an energy usually defined as the border between high energy (HE, $E>100\\,{\\rm MeV}$) and very high energy (VHE, $E>100\\,{\\rm GeV}$) where ACTs become sensitive. Of crucial importance for AGN studies evoked here is of course the LAT's uniform all-sky monitoring capabilities along with its efficiency, allowing observers at any other wavelength to have (almost) guaranteed GeV observations of their favourite sources without need to apply for observations.  Within a year, the LAT has collected $\\sim 200\\times 10^6$ $\\gamma$-rays, 2 orders of magnitude more than its predecessor. The first \\textit{Fermi} catalog\\cite{fermicat} (1FGL), based on 11 months of scientific operations, contains 1451 sources of which 1043 are at Galactic latitudes $b>10^\\circ$. From this latter sample was then derived the first LAT AGN catalog (1LAC), comprising 300 BL Lac objects (compared to 13 for EGRET), 296 flat-spectrum radio quasars (FSRQs, of which EGRET saw 66), and 113 AGN of other (or unknown) type. This is a tremendous change of the HE sky with just one of the 10 expected years of scientific operations, so even more is to be expected in the future.   \\def\\figsubcap#1{\\par\\noindent\\centering\\footnotesize(#1)} \\begin{figure}[t]% \\begin{center} \\parbox{2.1in}{\\epsfig{figure=fig1.eps,width=2.5in}\\figsubcap{a}} \\hspace*{1pt} \\parbox{2.1in}{\\epsfig{figure=fig2.eps,width=2.5in}\\figsubcap{b}} \\caption{(a) The \\textit{Fermi} blazar photon index as a function of its derived synchrotron peak frequency $\\nu_{\\rm peak}^{\\rm s}$, as derived in Ref.~\\refcite{fermi1lac}. (b) A similar study using the EGRET data, as published in Ref.~\\refcite{nandik}\\!. Thanks to a much larger population and a better energy resolution, a clear correlation appears where the blazar spectrum becomes increasingly harder when the synchrotron peak $\\nu_{\\rm peak}^{\\rm s}$ moves blueward from the optical to the X-rays.} \\label{fig1} \\end{center} \\end{figure}  \\subsection{Ground-based telescopes}  The most sensitive extensive air shower (EAS) detector to date is the Milagro observatory, which has provided a 1 to 100 TeV survey of the $\\gamma$-ray sky in which 34 \\textit{galactic} \\textit{Fermi} sources were found\\cite{milagro}\\!. So far however ACTs have been the dominant technique for studying the VHE extragalactic field, since the $\\gamma$-ray horizon at TeV energies is (at best) located at $z\\sim0.1$ due to photon-photon interactions with the extragalactic background light (EBL).   The new generation of major ACTs (HESS, MAGIC and VERITAS mainly), with larger telescopes and improved reconstruction techniques, have augmented the flux sensitivity of this technique by an order of magnitude within a decade, reaching a significance of 5 standard deviations on Crab nebula type sources in $\\sim 30\\,{\\rm s}$ or less, and ways are found to decrease the effect of moonlight in order to increase the observation time. The population of extragalactic ``TeV sources'' is currently at 30, but continuously increasing\\footnote{See the catalog TeVCat http://tevcat.in2p3.fr}\\!,as can be seen in Fig.~\\ref{fig2}. Most of the sources are BL Lacs with fluxes at the level of a few percent of the Crab nebula equivalent, with little measurable variability.   \\begin{figure}[t] \\begin{center} \\psfig{file=srctimeline.eps,width=5in} \\end{center} \\caption{The number of extragalactic sources detected by ACTs, as a funcion of their detection date, using the TeVCaT. The slope change from $\\sim 0.5/{\\rm year}$ to $\\sim 5/{\\rm year}$ coincides with the current generation of ACTs becoming operational, illustrating the dramatic effect of the sensitivity increase.} \\label{fig2} \\end{figure} ",
        "conclusions": ""
    },
    "1005/1005.0429_arXiv.txt": {
        "abstract": "The spin of a supermassive black hole (SMBH) is directly related to the radiative efficiency of accretion on to the hole, and therefore impacts the amount of fuel required for the black hole to reach a certain mass. Thus, a knowledge of the SMBH spin distribution and evolution is necessary to develop a comprehensive theory of the growth of SMBHs and their impact on galaxy formation. Currently, the only direct measurement of SMBH spin is through fitting the broad \\fe\\ line in AGNs. The evolution of spins could be determined by fitting the broad line in the integrated spectra of AGNs over different redshift intervals. The accuracy of these measurements will depend on the observed integrated line strength. Here, we present theoretical predictions of the integrated relativistic \\fe\\ line strength as a function of redshift and AGN luminosity. The equivalent widths of the integrated lines are much less than $300$~eV. Searches for the integrated line will be easiest for unobscured AGNs with $2$--$10$~\\kev\\ luminosities between $44 < \\log L_{X} \\leq 45$. The total integrated line makes up less than $4$\\% of the X-ray background, but its shape is sensitive to the average SMBH spin. By following these recommendations, future \\textit{International X-ray Observatory} surveys of broad \\fe\\ lines should be able to determine the spin evolution of SMBHs. ",
        "introduction": "\\label{sect:intro} In astrophysics, black holes can be described by only two parameters: their mass $M_{\\mathrm{BH}}$ and their angular momentum $J$ (parametrized by the dimensionless spin parameter $a \\equiv cJ/GM_{\\mathrm{BH}}^2$ where $0 \\leq a \\leq 1$). Over the last several decades, estimates of black hole masses have been obtained by measuring the gravitational influence of the hole on nearby objects such as binary stellar companions \\citep[e.g.,][]{bol72,cow83,mr06}, orbiting stars \\citep[e.g.,][]{ghez08}, stellar cusps \\citep[e.g.,][]{mag98} or dense clouds of ionized gas \\citep[e.g.,][]{peter04}. However, measurements of black hole spin are much more difficult to obtain as the spin only significantly affects the shape of space-time out to $\\approx 6$~$r_g$ ($r_g=GM_{\\mathrm{BH}}/c^2$ is a gravitational radius) from the event horizon \\citep{bpt72}. In particular, the value of the spin is directly related to the radius of the innermost stable circular orbit (ISCO) around the hole \\citep{bpt72}, and therefore impacts the radiative efficiency of thin accretion disks. If the accretion flow is rotating in the same sense as the hole, then a larger $a$ brings the ISCO closer to the event horizon, increasing the efficiency that an accretion flow can convert gravitational potential energy into radiation. A larger fraction of the rest-mass energy is radiated away in a highly efficient accretion disk, reducing the growth rate of the black hole mass. Therefore, measuring the distribution of spins will elucidate the role of steady accretion \\citep[e.g.,][]{v05}, chaotic accretion \\citep[e.g.,][]{kp06} and black hole mergers \\citep[e.g.,][]{bv08} in the build-up of supermassive black holes (SMBHs).  X-ray variability and optical microlensing data show that the X-ray source in AGNs is within 10~$r_g$ \\citep[e.g.,][]{iwa04,dai10}, indicating that the accretion disk is likely illuminated by X-rays down to the ISCO. As spin clearly impacts the radius of the ISCO, the extent of the red wing of the relativistically broadened \\fe\\ line is a direct measurement of black hole spin \\citep[e.g.,][]{rf08}. In recent years, long observations by \\xmm\\ and \\textit{Suzaku} have used this method to obtain SMBH spin estimates for a small number of bright, nearby Seyfert~1 galaxies \\citep{br06,min09,sch09}. These type of spin measurements depend on precise spectral fitting to the broad \\fe\\ line and continuum and are most easily performed on sources with strong \\fe\\ lines. Thus, the larger collecting area provided by the future \\textit{International X-ray Observatory} (\\ixo) will be necessary in order to extend these types of spin measurements over a significant range of redshift. In the meantime, attempts have been made to measure the average \\fe\\ line profile by stacking the X-ray data of the many AGNs detected in deep \\textit{Chandra} or \\xmm\\ surveys over a wide range of redshift and AGN luminosity. If a significant \\fe\\ signal can be detected in the integrated spectrum then, in principle, the average spin of the black holes contributing to the line could be determined. Again, the ability to make this measurement depends crucially on the strength of the integrated \\fe\\ line. \\citet{streb05} averaged the \\xmm\\ spectra of type 1 and type 2 AGNs detected in the Lockman Hole and found a strong (equivalent width [EW] $\\sim 500$~eV) and broad \\fe\\ line that indicated a non-zero spin for the average SMBHs in their sample. However, similar stacking analyses of different datasets by \\citet{bgc05} and \\citet{cor08} were not able to detect a broad component, with \\citet{cor08} finding that the EW of the broad component must be $< 400$~eV. These attempts to measure the integrated broad \\fe\\ line are technically challenging and are limited by the unknown intrinsic distribution of broad \\fe\\ line strengths.  Recently, \\citet{bal10} described a method by which the average EW of the broad \\fe\\ line could be calculated as a function of $z$ and luminosity. That paper concentrated on the distribution of likely EWs from individual sources, and did not consider the result of integrating lines over specific redshift intervals, a step crucial to measuring the spin evolution of SMBHs. This letter builds on the \\citet{bal10} method to calculate the integrated \\fe\\ line over several different redshift and luminosity ranges. We also consider the dependence of AGN obscuration on the appearance of the integrated \\fe\\ lines. These predictions show which regions in the redshift and AGN luminosity plane will have the largest average broad \\fe\\ EW, and will be vital for guiding current and future attempts to map out the spin evolution of SMBHs by fitting the average \\fe\\ line profile. The calculations are reviewed in the next section, and the results are presented in Sect.~\\ref{sect:res}. The following $\\Lambda$-dominated cosmology is assumed in this paper: $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.7$, and $\\Omega_{m}=0.3$ \\citep{spe03}. ",
        "conclusions": "\\label{concl} Measurements of the spin distribution of SMBHs will have profound consequences on our understanding of galaxy formation and evolution. The only direct way to measure SMBH spins is via the broad \\fe\\ line in AGNs, and attempts have been made to measure the \\fe\\ profile in the spectrum of AGNs that have been integrated over a range of $z$ and $\\log L_{X}$. The accuracy and precision of these measurements will be maximized if the spectral fitting can be performed on \\fe\\ lines with the largest possible EW. Therefore, this letter presented predictions of the integrated relativistic \\fe\\ EW as functions of $z$ and $\\log L_{X}$. Fig.~\\ref{fig:ews} showed that the integrated broad \\fe\\ EW will be significantly less than 300~eV for all ranges of $z$ and $\\log L_{X}$. In fact, if $R=1$ for most AGNs, the integrated broad component will have an EW between 100 and 200~eV. Taking into account the relative space density of AGNs, it is recommended that the search for integrated broad \\fe\\ lines be confined to AGNs with luminosities in the range $44 \\leq \\log L_{X} < 45$. It is further recommended that these experiments be limited to type 1 AGNs.  Stepping back from the quantitative results, this letter has shown that measuring the spin evolution of SMBHs can be done with the broad \\fe\\ line. In fact, \\ixo\\ will be perfectly placed to obtain this goal, as its sensitivity will allow it to obtain the high-quality spectra necessary to measure the broad \\fe\\ lines. \\ixo\\ spectral surveys of target lists provided by the \\chandra\\ deep fields will be a very powerful tool to trace the spin of SMBHs up to and beyond $z=1$."
    },
    "1005/1005.0790_arXiv.txt": {
        "abstract": "We consider generalizations of the entropic accelerating universe recently proposed in Ref.~\\cite{Easson:2010av,Easson:2010xf} and show that their background equations can be made equivalent to a model with a dark energy component with constant parameter of state $w_{X} = -1 + 2\\, \\gamma /3$, where $\\gamma$ is related to the coefficients of the new terms in the Friedman equations. After discussing all the Friedman equations for an arbitrary $\\gamma$, we show how to recover the standard scalings for dust and radiation. The acoustic scale $\\ell_A$, related to the peak positions in the pattern of the angular power spectrum of the Cosmic Microwave Background anisotropies, is also computed and yields the stringent bound $|\\gamma|\\ll 1$. We then argue that future data might be able to distinguish this model from pure $\\Lambda$CDM (corresponding to $\\gamma=0$). ",
        "introduction": "\\label{intro} Modern theoretical cosmology includes an early period of accelerated expansion named {\\em inflation\\/}~\\cite{inflation}, whose driving force is commonly modeled using a scalar field (the {\\em inflaton\\/}) of uncertain nature. A similarly accelerated phase is undergoing now~\\cite{supernovae} and has led to conceive the existence of an equally unspecified {\\em dark energy\\/} component in the matter content of the Universe \\cite{DarkEnergy}. \\par An alternative scenario has been recently proposed in Refs.~\\cite{Easson:2010av,Easson:2010xf}, based on the idea of entropic gravity introduced in Ref.~\\cite{verlinde}. In this context, the equations governing the time-evolution of the cosmic scale factor contain terms proportional to the Hubble function squared $H^2$ and its time derivative $\\dot H$ originating at the boundary of spatial sections of our universe. According to the authors of Refs.~\\cite{Easson:2010av,Easson:2010xf}, such terms could explain the acceleration occurring both in the early stages and at present. Boundary terms, whose nature is well-known in General Relativity~\\cite{Carroll:1997ar}, have indeed been analyzed in various contexts, for example in Refs.~\\cite{otherPaper}. \\par In this work, we will not analyze how these terms emerge from an action principle, nor if a unique Lagrangian can be defined at all. We shall instead assume general modifications of the form considered in Refs.~\\cite{Easson:2010av,Easson:2010xf} and then try to constrain their possible effectiveness by comparing the corresponding Cosmic Microwave Background (CMB) acoustic scale (see, e.g., Refs.~\\cite{Page:2003fa,Gruppuso:2005xy}) with the most recent available WMAP data~\\cite{Komatsu:2010fb}. Note that a standard Monte-Carlo Markov Chain analysis (usually employed to extract the cosmological parameters by comparison with available observations) is not feasible if the model is unknown at the linear order. On the contrary, the CMB acoustic scale can be computed directly from the background equations, and this will allow us to obtain a constraint for the free parameters of the model by comparing with the most recent CMB data. \\par The paper is organized as follows: in Section~\\ref{mfe} we obtain the complete set of Friedman equations for a generalization of the entropic models introduced in Refs.~\\cite{Easson:2010av,Easson:2010xf}. This, in Section~\\ref{effDE}, will allow us to regard the model in terms of an effective dark energy contribution depending on one parameter $\\gamma$. In particular, bounds on $\\gamma$ will be obtained by comparing the CMB acoustic scale with the 7yr~WMAP data in Section~\\ref{acousticscale}, after computing the deceleration parameter in Section~\\ref{decpar}. Conclusions will be drawn in Section~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions} In the present paper, we have shown how it is possible to recover standard background scalings for radiation and matter and standard effective cosmological equations [see Eqs.~(\\ref{EqFriedman_ddota_eff}) and (\\ref{EqFriedman_dota_eff})] when the Friedman equation for $\\ddot a$ is modified by adding terms proportional to $H^2$ and $\\dot H$ like in Eq.~(\\ref{eqFriedman}). An example that requires such a modification is given by the entropic accelerating universe of Refs.~\\cite{Easson:2010av,Easson:2010xf}, although our considerations are more general. Moreover we have shown how to obtain the recent cosmological acceleration within the considered model, without adding a dark energy fluid. We note that for the range of parameters considered here, the model under analysis does not modify the evolution of the universe when it was matter or radiation dominated. Therefore, none of the standard cosmological constraints coming from such early epochs, as for instance the Big Bang Nucleosynthesis (BBN), are affected. \\par Specifically, we have shown that the parameter space admits a region (i.e., $\\gamma < 1$) where the universe accelerates at recent cosmological times (i.e., $z \\sim 0.5$). In fact, the additional terms mimic the behavior of a fluid with a constant parameter of state $w_X = -1 + 2\\, \\gamma /3$. This has been studied by computing the deceleration parameter $q$ [see Fig.~\\ref{FigDue} and Section~\\ref{decpar}] and stringent constraints have been obtained comparing the CMB acoustic scale $\\ell_A$ with the WMAP~7yr release data. Note that a standard Monte-Carlo Markov Chain analysis (usually employed to extract the cosmological parameters by comparison with available observations) is not feasible if the model is not known at linear order~\\footnote{In fact, as far as we know, no Lagrangian is known for these models. However, as we stated in the Introduction, we are not debating the theoretical ground the ``entropic-like'' proposal is based on, but are rather interested in which constraints we can provide for such class of models from what we have at hand, i.e.~the background equations.}. On the contrary, the CMB acoustic scale can be computed directly from the background equations, and this has allowed us to obtain a constraint for the free parameters of the model by comparing with the most recent CMB data. This comparison has told us that $|\\gamma|\\ll 1$ (so that $w_X\\simeq -1$) and the coefficients of $\\dot H$ and of $H^2$ in Eq.~(\\ref{eqFriedman}) must therefore satisfy Eq.~(\\ref{cCC}) for the model to be phenomenologically viable. In particular, the entropic accelerating universe corresponds to a specific choice of the constants $C_H$ and $C_{\\dot H}$, that is $\\gamma_{\\rm I} = 0$ and $\\gamma_{\\rm II}=0.68$ for the two cases explicitly mentioned in Ref.~\\cite{Easson:2010av}. The latter is at odd with the constraint~(\\ref{rangepergamma}), whereas the former is consistent. \\par Future CMB observations coming from the {\\sc Planck} satellite~\\footnote{Planck (http://www.esa.int/Planck) is a project of the European Space Agency, ESA.} are expected to improve the error on the acoustic scale by about one order of magnitude~\\cite{Colombo:2008ta}. The same improvement is therefore expected for the estimate of the parameter $\\gamma$, which, in principle, should allow us to distinguish these models from the pure $\\Lambda$CDM model with $\\gamma=0$. We finally mention that our findings are in agreement with those of Ref.~\\cite{koivisto}."
    },
    "1005/1005.5456_arXiv.txt": {
        "abstract": "\\noindent Dwarf and low surface brightness galaxies are ideal objects to test modified Newtonian dynamics (MOND), because in most of these galaxies the accelerations fall below the threshold below where MOND supposedly applies. We have selected from the literature a sample of 27 dwarf and low surface brightness galaxies. MOND is successful in explaining the general shape of the observed rotation curves for roughly three quarters of the galaxies in the sample presented here. However, for the remaining quarter, MOND does not adequately explain the observed rotation curves.  Considering the uncertainties in distances and inclinations for the galaxies in our sample, a small fraction of poor MOND predictions is expected and is not necessarily a problem for MOND.  We have also made fits taking the MOND acceleration constant, $a_0$, as a free parameter in order to identify any systematic trends.  We find that there appears to be a correlation between central surface brightness and the best-fit value of $a_0$, in the sense that lower surface brightness galaxies tend to have lower $a_0$. However, this correlation depends strongly on a small number of galaxies whose rotation curves might be uncertain due to either bars or warps. Without these galaxies, there is less evidence of a trend, but the average value we find for $a_0\\approx 0.7 \\times 10^{-8}$ cm s$^{-2}$ is somewhat lower than derived from previous studies. Such lower fitted values of $a_0$ could occur if external gravitational fields are important. ",
        "introduction": "\\noindent Modified Newtonian dynamics (MOND) was proposed by Milgrom (1983a,b) as an alternative to dark matter.  MOND posits that the effective gravitational force deviates from the Newtonian force: at accelerations below a critical value, $a_0$, the gravitational force is proportional to the square root of the Newtonian force; at higher accelerations, the force is Newtonian.  In principle, MOND predicts the shape and amplitude of the observed rotation curves from the observed mass distribution, i.e., gas and stars, with only the mass-to-light ratio (M/L) of the stellar disk as adjustable parameter.  The rotation curves of a large number of spiral galaxies have been considered in the context of MOND. For example, Begeman et al.\\ (1991), Sanders (1996), Sanders \\& Verheijen (1998), and Sanders \\& Noordermeer (2007) demonstrated that for the spiral galaxies the observed rotation curves were, in most cases, predicted in detail, using the MOND prescription, from the observed light and gas distributions.  Moreover, the M/Ls derived from these fits are generally astrophysically plausible and consistent with stellar population synthesis models.  In a handful of these galaxies the rotation curves predicted by MOND are noticeably different from the observed rotation curves. These were not considered falsifications of MOND because for most of these galaxies the authors identified an obvious problem with the observed rotation curve (e.g., uncertain inclination or distance) or with the use of the rotation curve as a tracer of the gravitational force (e.g., disturbed velocity fields).  Although MOND appears to explain the magnitude of the discrepancy in spiral galaxies, it has has long been known that the theory predicts more matter than is observed in clusters of galaxies (e.g., Sanders 1999).  In this sense the famous \"bullet cluster\", which has been presented as direct emprical evidence for the existence of dark matter (e.g., Clowe et al.\\ 2006), does not pose a new problem for MOND. Additional undetected matter in some dissipationless form does seem to be required in clusters in the context of MOND.  In any case, additional required matter is not formally a falsification of MOND and in no sense detracts from the success of this algorithm on the scale of galaxies.  \\begin{deluxetable*}{lcrrrrrrrr}[th] \\tabletypesize{\\scriptsize} \\tablecaption{Galaxy properties\\label{tabsample}} \\tablewidth{0pt} \\tablehead{ \\colhead{Name} & \\colhead{Source} & \\colhead{$\\mathrm{D}_a$} & \\colhead{$\\mu_0^R$} & \\colhead{$h$} & \\colhead{$M_R$} & \\colhead{$M_\\mathrm{HI}$} & \\colhead{$i$} & \\colhead{$v_\\mathrm{rot}$} & \\colhead{$a_\\mathrm{lmp}/a_0$} \\\\ \\colhead{} & \\colhead{} & \\colhead{Mpc} & \\colhead{mag$ \\prime\\prime^{-2}$} & \\colhead{kpc} & \\colhead{mag} & \\colhead{$10^8 M_\\odot$} & \\colhead{$^\\circ$} & \\colhead{km s$^{-1}$} & \\colhead{} \\\\ \\colhead{(1)} & \\colhead{(2)} & \\colhead{(3)} & \\colhead{(4)} & \\colhead{(5)} & \\colhead{(6)} & \\colhead{(7)} & \\colhead{(8)} & \\colhead{(9)} & \\colhead{(10)} \\\\ } \\startdata UGC 731 & ~SMBB ~& 8.0 &23.0 &1.65 &-16.6 &7.37 &57 &  74 & 0.25  \\\\ UGC 3371 & ~SVBA ~&12.8 &23.3 &3.09 &-17.7 &12.2 &49 &  86 & 0.23  \\\\ UGC 4173 & ~dBB ~&18.0 &24.3 &4.77 &-18.0 &24.3 &40 &  57 & 0.08  \\\\ UGC 4325 & ~SMBB ~&10.8 &21.6 &1.74 &-18.2 &8.61 &41 &  93 & 0.44  \\\\ UGC 4499 & ~SMBB ~&13.9 &21.5 &1.59 &-17.9 &13.6 &50 &  74 & 0.20  \\\\ UGC 5005 & ~dBB ~&  56 &22.9 &4.71 &-18.8 &33.2 &41 &  99 & 0.11  \\\\ UGC 5414 & ~SSAH ~&10.7 &21.8 &1.59 &-17.7 &7.42 &55 &  61 & 0.26  \\\\ UGC 5721 & ~SMBB ~& 6.6 &20.2 &0.45 &-16.6 &6.63 &61 &  79 & 0.28  \\\\ UGC 5750 & ~dBMR ~&  60 &22.6 &5.99 &-19.6 &11.4 &64 &  79 & 0.09  \\\\ UGC 6446 & ~SSAH ~&12.8 &21.4 &2.00 &-18.5 &15.4 &52 &  80 & 0.20  \\\\ UGC 7232 & ~SSAH ~& 3.5 &20.2 &0.33 &-15.3 &0.71 &59 &  44 & 0.62  \\\\ UGC 7323 & ~SSAH ~& 8.7 &21.2 &2.35 &-19.0 &8.48 &47 &  86 & 0.38  \\\\ UGC 7399 & ~SSAH ~& 8.4 &20.7 &0.79 &-17.1 &7.40 &55 & 109 & 0.35  \\\\ UGC 7524 & ~SSAH ~& 4.6\\rlap{$^\\ast$} &22.2 &3.37 &-18.7 &16.7 &46 &  84 & 0.22  \\\\ UGC 7559 & ~SSAH ~& 4.9\\rlap{$^\\ast$} &23.8 &1.02 &-14.6 &1.71 &61 &  33 & 0.11  \\\\ UGC 7577 & ~SSAH ~& 2.5\\rlap{$^{\\ast\\ast}$} &22.5 &0.60 &-14.9 &0.42 &63 &  18 & 0.06  \\\\ UGC 7603 & ~dBB ~& 7.3 &20.8 &0.96 &-17.0 &6.13 &78 &  64 & 0.21  \\\\ UGC 8490 & ~SMBB ~& 4.6\\rlap{$^\\ast$} &20.5 &0.62 &-17.1 &7.01 &50 &  80 & 0.21  \\\\ UGC 9211 & ~dBB ~&12.6 &22.6 &1.32 &-16.2 &10.5 &44 &  66 & 0.17  \\\\ UGC 11707 & ~SMBB ~&17.0 &23.1 &4.61 &-18.8 &42.7 &68 & 100 & 0.20  \\\\ UGC 11861 & ~SMBB ~&  27 &21.4 &6.48 &-20.9 &81.7 &50 & 153 & 0.43  \\\\ UGC 12060 & ~SSAH ~&16.8 &21.6 &1.89 &-18.1 &20.7 &40 &  75 & 0.16  \\\\ UGC 12632 & ~SSAH ~& 6.9 &23.5 &2.57 &-17.1 &8.66 &46 &  76 & 0.22  \\\\ F568-V1 & ~SMBB ~&  86 &22.8 &3.42 &-18.8 &28.1 &40 & 118 & 0.24  \\\\ F574-1 & ~SMBB ~& 103 &22.6 &4.60 &-19.3 &39.7 &65 & 100 & 0.22  \\\\ F583-1 & ~dBMR ~&  34 &22.0 &1.71 &-17.5 &19.9 &63 &  84 & 0.15  \\\\ F583-4 & ~dBMR ~&  52 &22.4 &2.89 &-17.8 &6.30 &55 &  70 & 0.22  \\\\ \\enddata \\tablecomments{(1) the name of the galaxy, (2) source of the rotation curve: de Blok \\etal\\ (2001, dBMR), de Blok \\& Bosma (2002; dBB), Swaters \\etal\\ (2003a; SMBB), Swaters \\etal\\ (2003b; SVBA), Swaters \\etal\\ (2009; SSAH), (3) adopted distance, taken from Swaters \\& Balcells (2002), except that distances for galaxies marked with $^\\ast$ were taken from Karachentev et al.\\ (2003), those marked with $^{\\ast\\ast}$ were taken from M\\'endez et al. (2002) (4--6) surface brightness, scale length, and absolute magnitude, from Swaters \\& Balcells (2002), de Blok et al.\\ (1995), and McGaugh \\& Bothun (1994), (7) HI mass, from Swaters et al.\\ (2003), de Blok et al.\\ (1996), and van der Hulst et al.\\ (1993) (8) inclinations, from Swaters et al.\\ (2004) and de Blok et al.\\ (1996), (9) maximum rotation velocity, (10) acceleration at the last measured point of the rotation curve.} \\end{deluxetable*}  MOND is more strongly tested by spiral galaxies with well-determined distances because the $a_0$ depends inversely upon distance.  Bottema et al.\\ (2002) considered a sample of 4 spiral galaxies with well-measured rotation curves and Cepheid-based distances determined as part of the HST program on the extragalactic distance scale (Sakai et al.\\ 2000). Two galaxies are in good agreement with MOND predictions, but for NGC~2841 and NGC~3198 the rotation curves are not in agreement with MOND, unless the distance to NGC~2841 is at least 20\\% larger, and the distance to NGC~3198 is at least 10\\% smaller than their Cepheid distances indicate. Bottema et al.\\ (2002) argued that NGC~3198 is probably consistent with MOND within the uncertainties of the Cepheid distance. With respect to NGC 2841, Milgrom and Sanders (2008) have demonstrated that a reasonable match to the observed rotation curve can be achieved with alternative forms of the function that described the transition from the Newtonian to the MOND regime.  Dwarf and low surface brightness (LSB) galaxies provide an especially good test for MOND, because in most of these galaxies the acceleration is below the MOND threshold of $a_0 \\approx 10^{-8}$~cm~s$^{-2}$ (determined empirically from rotation curve fits, see e.g., Sanders \\& McGaugh 2002) at all radii, whereas in high surface brightness spiral galaxies the acceleration is usually above the threshold in the inner parts.  Milgrom (1983b) predicted that LSB galaxies should exhibit a large discrepancy between the detectable and Newtonian dynamical mass within the optical disk, and this prediction has since been confirmed (e.g., de Blok \\& McGaugh 1997; McGaugh \\& de Blok 1998; Swaters \\etal\\ 2000).  Apart from this general prediction, the MOND prescription is successful in predicting the detailed rotation curves of LSB galaxies as well. De Blok \\& McGaugh (1998) found that for their sample of 15 LSB galaxies all rotation curves are well fitted with MOND, although for 6 galaxies modest adjustments of the inclination were necessary in order to get an acceptable match to the rotation curve.  The sample presented in Sanders \\& Verheijen (1998) contains 12 LSB galaxies, and most of these the rotation curves can be well explained by MOND as well.  \\begin{figure*}[p] \\resizebox{0.9\\hsize}{!}{\\includegraphics{figure1.eps}} \\caption{MOND predictions. The dots represent the measured rotation velocities and their errors. The dashed line represents the contribution of the HI to the rotation curve. In cases where there is little or no HI at small radii in the HI disks, the HI at larger radii produces a net outward gravitational force, which is represented here by the negative velocities. The thin solid line represents the contribution of the stars. The thick solid line is the best MOND fit with only the stellar mass-to-light ratio as a free parameter. The dotted line is the best MOND fit with both the mass-to-light ratio and the distance as free parameter. \\label{figfits}} \\end{figure*}  The results for dwarf galaxies show a fairly large spread.  Milgrom \\& Braun (1988) concluded that the observed rotation curve of DDO~154 is in very good agreement with the rotation curve predicted by MOND. Lake (1989) considered a sample of 6 dwarf galaxies in the context of MOND and concluded that acceptable fits require systematically lower values for $a_0$ than were found for spiral galaxies.  In addition, Lake (1989) found an apparent correlation between the asymptotic rotation velocity and the fitted value of $a_0$.  Milgrom (1991), however, pointed out that uncertainties in inclinations and distances can well explain the fact that MOND appears to fail for these galaxies. The rotation curves of the dwarf galaxies included in the samples of Begeman et al.\\ (1991) and Sanders (1996) also agree well with those calculated with MOND. On the other hand, Blais-Ouellette et al.  (2001) reported that MOND provides poor predictions of the rotation curves of the three dwarf galaxies in their sample, and reported that a higher value of $a_0$ is needed to make their rotation curves derived from Fabry-Perot observations consistent with MOND. More recently, Milgrom \\& Sanders (2007) concluded that the rotation curves of the four low-mass galaxies in their sample are explain well in the context of MOND.  Given this apparent spread in conclusions on whether MOND correctly predicts the rotation curves of dwarf galaxies, we here test MOND against a sample of rotation curves of dwarf and LSB galaxies that have recently become available in the literature. We included both dwarf and LSB galaxies because the physical properties of these galaxies are similar, although some dwarf galaxies have high surface brightness, and some LSB galaxies are large in size.  There are several difficulties inherent to such a sample. For example, because most dwarf galaxies are nearby, different methods to estimate the distance have been used, and this may introduce a dispersion in the derived parameters.  Inclination is another uncertainty for a class of objects which are known to be irregular with, in some cases, large scale asymmetries.  Moreover, the assumption that the rotation curve is a tracer of the radial force distribution may, itself, be questionable in cases where asymmetries are present.  Bearing these considerations in mind, one might not expect, a priori, the agreement between the predicted and observed rotation curves to be as good as noted for earlier samples. A systematic failure, however, would be problematic for MOND.  The outline of this paper is as follows.  In Section~\\ref{secsample} we describe the sample used in this paper. In Section~\\ref{secfitproc} we describe the fitting procedure, and in Section~\\ref{secfitres} the fitting results. Next, we comment on individual galaxies in Section~\\ref{secindiv}.  We describe the uncertainties in Section~\\ref{secuncert}, and we discuss the results in Section~\\ref{secdisc}. Finally, we present our conclusions in Section~\\ref{secconc}. ",
        "conclusions": "\\label{secconc}  \\noindent We have presented MOND fits for a sample of 27 dwarf and LSB galaxies.  MOND is remarkably successful at predicting the general shape of the rotation curves in this sample. Only for approximately a quarter of the galaxies MOND does not adequately predict the observed rotation curves.  This is remarkable given the uncertainties associated with the galaxies in this sample, especially the uncertainties in distance and inclination. Close inspection of the fits indicates that for almost 80\\% of the galaxies MOND predicts higher velocities in the outer rotation curves than is observed. Although this could be the result of uncertainties, this systematic deviation can also be explained if the MOND acceleration parameter is slightly lower than usually assumed, with a value of $a_0=2150$ km$^2$ s$^{-2}$ kpc$^{-1}$ (or $0.7 \\times 10^{-8}$ cm~s$^{-2}$). We find that there appears to be some evidence of a correlation between $a_0$ and the central extrapolated disk surface brightness $\\mu_R$, but this possible correlation depends heavily on a few galaxies at high surface brightness, whose rotation curves may be uncertain because of bars or warps.  At the low surface brightness end, the derived values of $a_0$ could be be biased towards low values if external fields are important.  Overall, the uncertainties associated with this sample make it difficult to draw strong conclusions. Improved distances and inclination estimates for these galaxies could make it possible to test MOND more strongly."
    },
    "1005/1005.4982_arXiv.txt": {
        "abstract": "We present local two-dimensional (2D) and three-dimensional (3D) hybrid numerical simulations of particles and gas in the midplane of protoplanetary disks (PPDs) using the Athena code. The particles are coupled to gas aerodynamically, with particle-to-gas feedback included. Magnetorotational turbulence is ignored as an approximation for the dead zone of PPDs, and we ignore particle self-gravity to study the precursor of planetesimal formation. Our simulations include a wide size distribution of particles, ranging from strongly coupled particles with dimensionless stopping time $\\tau_s\\equiv\\Omega t_{\\rm stop}=10^{-4}$ to marginally coupled ones with $\\tau_s=1$ (where $\\Omega$ is the orbital frequency, $t_{\\rm stop}$ is the particle friction time), and a wide range of solid abundances. Our main results are: 1. Particles with $\\tau_s\\gtrsim10^{-2}$ actively participate in the streaming instability, generate turbulence and maintain the height of the particle layer before Kelvin-Helmholtz instability is triggered. 2. Strong particle clumping as a consequence of the streaming instability occurs when a substantial fraction of the solids are large ($\\tau_s\\gtrsim10^{-2}$) and when height-integrated solid to gas mass ratio $Z$ is super-solar. We construct a toy model to offer an explanation. 3. The radial drift velocity is reduced relative to the conventional Nakagawa-Sekiya-Hayashi (NSH) model, especially at high $Z$. Small particles may drift outward. We derive a generalized NSH equilibrium solution for multiple particle species which fits our results very well. 4. Collision velocity between particles with $\\tau_s\\gtrsim10^{-2}$ is dominated by differential radial drift, and is strongly reduced at larger Z. This is also captured by the multi-species NSH solution. Various implications for planetesimal formation are discussed. In particular, we show there exist two positive feedback loops with respect to the enrichment of local disk solid abundance and grain growth. All these  effects promote planetesimal formation. ",
        "introduction": " ",
        "conclusions": "\\subsection[]{Summary of Main Results}  The main purpose of this paper is to study the dynamics of solids and gas in the midplane of PPDs using hybrid simulations. The solids and gas are coupled aerodynamically, characterized by the dimensionless stopping time $\\tau_s=\\Omega t_{\\rm stop}$. We consider a wide size distribution of solids as an approximation to the outcome of grain growth in PPDs, ranging from sub-millimeter to meter size. The key ingredient of our simulations is the inclusion of feedback from particles to gas. Feedback is important when the local particle to gas mass ratio exceeds order unity. Moreover, it is essential for the generation of SI and KHI. In our simulations, we assume no external source of turbulence, as an approximation for the dead zone of PPDs. Turbulence in the disk midplane is generated self-consistently from the SI (driven by the radial pressure gradient in the gas) and/or KHI (driven by vertical shear). Our simulations are local, since very high numerical resolution is essential to resolve the SI and KHI. Self-gravity is ignored, as we focus on the particle-gas dynamics before the formation of planetesimals.  Our simulations are characterized by three sets of dimensionless parameters, namely the particle size distribution $\\tau_k$, solid abundance $Z$, and a parameter $\\Pi$ characterizing the radial pressure gradient. In this paper, we fix $\\Pi=0.05$, as appropriate for a wide range of disk model parameters (see \\S\\ref{ssec:scaling}). The dependence of the particle clumping properties on $\\Pi$ is presented in a separate paper \\citep{BaiStone10c}. We consider a flat mass distribution in logarithmic bins in $\\tau_s$, and vary $Z$ from $0.01$ to $0.03$ (see Table \\ref{tab:simulation}). We conduct both 2D and 3D simulations, where 2D simulations are performed in the radial-vertical plane in order for the SI to be actively generated. We run the simulations for $40-200$ orbits and study the properties of the particles and gas in the saturated state. The main results are summarized below.  \\begin{enumerate} \\item SI plays the dominant role in the dynamics of PPD midplane when the largest solids have stopping times $\\tau_s\\gtrsim10^{-2}$. Particles with $\\tau_s\\gtrsim10^{-2}$ actively participate in SI, while smaller particles behave passively. KHI is not observed in all our simulations, which suggests that it may be important only when all particles have $\\tau_s\\lesssim10^{-2}$.  \\item The strength of the turbulence generated by the SI and the scale height of the particle layer are self-regulated. There exists some threshold solid abundance, above which increasing $Z$ will result in weaker turbulence, which promotes particle settling, leading to rapid drop of the thickness of the particle layer and strong particle clumping.  \\item SI can concentrate particles into dense clumps with solid density exceeding the Roche density, which acts as the prelude of planetesimal formation. The particle clumping generally requires the presence of relatively large particles with $\\tau_s\\gtrsim10^{-2}$. It also sensitively depends on solid abundance, in favor of super-solar metallicity.  \\item The dense particle clumps are mostly made of the largest particles with size range spanning less than one order of magnitude. These particles are trapped in the clumps for several orbital times before leaving the clumps, providing a way for large particles to survive radial drift.  \\item The mean radial drift velocity for each particle species agrees well with a multi-species NSH equilibrium solution (see Appendix A). Strongly coupled particles drift outward, and the radial drift velocity for particles with larger $\\tau_s$ is strongly reduced relative to the conventional single-species NSH value, especially at large $Z$. This can increase the lifetime of the largest particles by a factor of a few.  \\item Turbulence generated by the SI leads to radial diffusion of particles, but the diffusion is slow and its effect is negligible compared with radial drift after about $100$ orbital periods (for the largest particles). Particle clumping effectively enhances radial diffusion by retaining a fraction of large particles in the clumps.  \\item Mutual collision velocity between $\\tau_s\\gtrsim10^{-2}$ particles is dominated by the difference in their radial drift velocities, and agrees well with calculations using the multi-species NSH equilibrium. The collision velocity is strongly reduced towards large disk metallicity relative to predictions from single-species NSH solution. Collision velocity induced by SI turbulence is only secondary.  \\end{enumerate}  \\subsection[]{Implications for Planetesimal Formation}  In this subsection we combine the results summarized in the previous subsection and discuss various implications for planetesimal formation. In particular, we emphasize that the importance of the local enrichment of solid materials in PPDs on planetesimal formation (two feedback loops, see \\S\\ref{ssec:loop1} and \\S\\ref{ssec:loop2}). Our logical chain is summarized in Figure \\ref{fig:summary}, and we elaborate various aspects of this diagram in the following.   \\begin{figure} \\centering \\includegraphics[width=80mm,height=90mm]{summary.eps} \\caption{Summary of logical relations between various factors relevant to planetesimal formation. A solid arrow from A to B indicates that A promotes B. Dashed double arrow indicates speculated connection between the two processes. $Z^{\\rm large}$ represents metallicity for large particles (millimeter or larger), $Z^{\\rm small}$ represents the abundance for micron-sized small grains. Radial pile-up denotes the pile-up of particles due to the positive radial gradient of the radial drift velocity. We refer to ``Grain growth\" as collisional coagulation towards the largest solid size. ``Dead zone\" refers to the radial extent of the dead zone. See text for details. }\\label{fig:summary} \\end{figure}   \\subsubsection[]{Conditions for Strong Particle Clumping}  Our simulations show a dichotomy in the parameter space in which strong particle clumping occurs. Strong clumping requires the presence of relatively large particles with $\\tau_s\\gtrsim10^{-2}$, and the abundance of these particles to be super-solar. These two requirements are represented by the two arrows connecting $Z^{\\rm large}$ and ``grain growth\" to ``planetesimals\" in Figure \\ref{fig:summary}. These two requirements can compensate each other: to form planetesimals, less grain growth is required if the solid abundance is large enough. In the case that all solids are strongly coupled to the gas, enhancing the disk metallicity by a factor of 5-30 may be able to trigger GI followed by planetesimal formation \\citep{Sekiya98,YoudinShu02,Chiang08}.   The particle size that corresponds to $\\tau_s=10^{-2}$ depends on the disk model and distance from the central star. In the MMSN model, this stopping time corresponds to $2$cm particles at 1AU, and $2$mm or smaller particles at 5AU or beyond, according to equation (\\ref{eq:taus}). These particle sizes are about the maximum particle size obtained in the recent dust coagulation calculations \\citep{Zsom_etal10}. For more massive disks, the corresponding particle size will be enlarged by a factor of $f_g$ in the Epstein regime. Since grain growth becomes more difficult when the particle size exceeds a millimeter, the SI scenario of planetesimal formation prefers less massive disks, or the outer part of the disk.  \\subsubsection[]{Enrichment of Local Solid Abundance}\\label{ssec:loop1}  Enrichment of the local abundance of solids is possible by several effects, already briefly mentioned in \\S\\ref{sec:intro}. Here we focus on the mechanism proposed by \\citet{YoudinShu02}. For a particle at a fixed size (small, with $\\tau_s<1$ at all disk radii considered), the (single-species) NSH radial drift velocity in PPDs decreases as the particle drifts inward. Therefore, radial drift causes the particles to pile-up towards the inner regions, leading to enhancement of the local abundance of solids. Depending on the time scale, the largest enhancement factor can reach $3-10$ in $10^5-10^6$ years \\citep{YoudinShu02,YoudinChiang04}, starting in the outer disk and moving inwards. This process corresponds to the arrow pointing from  ``radial pile-up\" to $Z^{\\rm large}$ in Figure \\ref{fig:summary}.  In this paper, we have shown that the radial drift velocity is reduced when local abundance of solids increases. This effect provides positive feedback to the enrichment process: particles pile up because the radial drift velocity is smaller at smaller disk radii. The enrichment of solid material at small disk radii further reduces the radial drift velocity, leading to stronger pile-ups. This effect corresponds to the arrow pointing from $Z^{\\rm large}$ to ``radial pile-up\" in Figure \\ref{fig:summary}. We see that the enrichment of the abundance of solids and the particle pile-up form a feedback loop that enhance each other. Therefore, we expect even stronger solid enrichment in the inner region of PPDs than previous calculations \\citep{YoudinShu02,YoudinChiang04} sufficient for SI and/or GI to form planetesimals within PPD dead zone.  \\subsubsection[]{Implications for Grain Growth}\\label{ssec:loop2}  The radial drift velocity adopted in grain coagulation models (e.g., \\citealp{Brauer_etal08a,Birnstiel_etal10,Zsom_etal10}) is generally taken from a single-species NSH equilibrium. However, we have shown that as particles settle to disk midplane, the radial drift velocity is reduced due to multi-species effects, and smallest particles can even drift outward rather than inward. More sophisticated modeling of grain growth is needed to incorporate the multi-species effects.  One consequence of the multi-species NSH equilibrium is that the enhancement of local abundance of solids strongly reduces the radial drift velocity, hence the particle mutual collision velocity. Because the radial drift velocity dominates the collision velocities for relatively large particles, we expect particles to grow larger in regions with large abundance of solids, due to the reduced collision velocity there. This effect is illustrated as an arrow pointing from $Z^{\\rm large}$ to ``grain growth\" in Figure \\ref{fig:summary}. In turn, grain growth form large particles from smaller ones, increasing the population of large particles, hence $Z^{\\rm large}$, leading to the arrow pointing in the opposite direction. Again, these two effects form a feedback loop, promoting grain growth and enrichment of solid material.   \\subsubsection[]{Influence from Magnetic Activity}  All our simulations ignore external sources of turbulence, particularly the MRI, by working in the dead zone of PPDs. Even very weak external turbulence with $\\alpha\\sim10^{-4}$ may stir up the solids and maintain them at a height that may be insufficient for SI to form dense particle clumps (see Table \\ref{tab:vertdiff}). \\citet{Johansen_etal07} showed that planetesimal formation is facilitated by MRI due to the long-lived overdensity regions that effectively trap the particles. However, they chose very large particle size with $\\tau_s=0.25-1$. When smaller particles are used, they are diffused to a much larger height. Although SI may still be present with MRI turbulence and form elongated structures \\citep{Balsara_etal09}, particle overdensities are small. Moreover, the mutual particle collision velocity is much higher in MRI turbulence \\citep{Carballido_etal08}, which inhibits particle growth. Therefore, we expect the dead zone to be the more favored region for planetesimal formation, and planetesimal formation should be easier in PPDs with a larger (radial) extent of the dead zone. This is shown as the arrow pointing from ``dead zone\" towards ``planetesimal formation\" in Figure \\ref{fig:summary}.  At the same time, the vertical and radial extent of the dead zone in PPDs strongly depend on the abundance of micron-sized and smaller grains \\citep{IlgnerNelson06,BaiGoodman09}, which is reflected in the arrow connecting $Z^{\\rm small}$ and ``dead zone\". Whether there is any connection between $Z^{\\rm large}$ and $Z^{\\rm small}$ is uncertain. Both $Z^{\\rm large}$ and $Z^{\\rm small}$ should be correlated with the overall disk metallicity. Moreover, with more large particles (larger $Z^{\\rm large}$, collisional fragmentation may lead to more small grains (larger $Z^{\\rm small}$). Despite many uncertainties, we draw a dashed double-arrow between $Z^{\\rm large}$ and $Z^{\\rm small}$ as our speculated connection between the two particle populations. If such connection exist, it represents the third way for (local) solid abundance enrichment to promote planetesimal formation.   \\subsection[]{Limitations and Outlook}  Our simulations take a local patch from a simple global disk model in which all physical quantities follow a power-law dependence on disk radii. The global structure of PPDs may be more complicated. In particular, the presence of a dead zone in PPDs changes the steady-state disk surface density profile, and may lead to local pressure maxima at the snow line \\citep{KretkeLin07,Brauer_etal08b}, and the inner edge of the dead zone \\citep{Dzyurkevich_etal10}. These pressure bumps are able to trap particles very efficiently. Moreover, solids can also be trapped in long-lived vortices \\citep{BargeSommeria95,KlahrHenning97, KlahrBodenheimer03,Johansen_etal04,Lyra_etal09}, although their existence is in debate. Finally, the structure of PPDs may be very non-steady, and undergo periods of outburst \\citep{Zhu_etal09a}. Global models take into account large-scale variations of disk structure and can follow the disk evolution. Currently, the numerical resolution required for resolving the SI is prohibitively high for running 3D global simulations. Nevertheless, one can perform local simulations in different disk environments, and piece them together to form a global picture, as in \\citet{Tilley_etal10}.  Our simulations focus on the dynamics in the vicinity of the disk midplane with very limited radial and vertical box sizes and no magnetic field. This is mainly constrained by the fine grid resolution required for this study.  In reality, MRI turbulence in the active layers may excite vertical oscillations in the disk midplane\\citep{FlemingStone03,OshiMacLow09}. Moreover, turbulent diffusion of ions may produce magnetic activity at the disk midplane, making it ``undead\" \\citep{TurnerSano08}. Finally, turbulent mixing of particles may become important when the active layer is relatively thick \\citep{Turner_etal10}. Including these effects into numerical simulations of planetesimal formation requires enlarging our box size in all dimensions by a factor of at least 10. Moreover, 3D rather than 2D simulations are necessary to maintain sustained MRI turbulence.  Such simulations are computationally expensive, however recently the static mesh refinement (SMR) algorithm in the Athena MHD code has been parallelized.  The cost of hybrid (particles and gas) simulations of layered disks can be substantially reduced by using fine resolution at the disk midplane (to capture the SI) and in the active layers (to capture the MRI), while using coarse resolution everywhere else. With SMR, it becomes feasible to study the effect of (non-ideal) MRI turbulence in the active layer on particle dynamics and planetesimal formation in the disk midplane. This is planned as future work."
    },
    "1005/1005.3270_arXiv.txt": {
        "abstract": "Off-equatorial circular orbits with constant latitudes (halo orbits) of electrically charged particles exist near compact objects. In the previous paper, we discussed this kind of motion and demonstrated the existence of minima of the two-dimensional effective potential which correspond to the stable halo orbits.  Here, we relax previous assumptions of the pseudo-Newtonian approach for the gravitational field of the central body and study properties of the halo orbits in detail. Within the general relativistic approach, we carry out our calculations in two cases. Firstly, we examine the case of a rotating magnetic compact star. Assuming that the magnetic field axis and the rotation axis are aligned with each other, we study the orientation of motion along the stable halo orbits. In the poloidal plane, we also discuss shapes of the related effective potential halo lobes where the general off-equatorial motion can be bound. Then we focus on the halo orbits near a Kerr black hole immersed in an asymptotically uniform magnetic field of external origin.  We demonstrate that, in both the cases considered, the lobes exhibit two different regimes, namely, one where completely disjoint lobes occur symmetrically above and below the equatorial plane, and another where the lobes are joined across the plane. A possible application of the model concerns the structure of putative circumpulsar discs consisting of dust particles. We suggest that the particles can acquire a small (but non-zero) net electric charge, and this drives them to form the halo lobes. ",
        "introduction": "Introduction} Motion of matter near compact stars and black holes has been discussed thoroughly in textbooks~\\cite{Chandrasekhar,Mis-Tho-Whe:1973}.  In the test-particle  approximation, the hydrodynamical terms are neglected, as expected in circumstances when the medium is rarefied and the mean free path is comparable with the typical length-scale of the system~\\cite{Pra:1980,Bic-Stu-Bal:1989,Stu-Bic-Bal:1999}. The interplay between gravitational attraction and the electromagnetic attraction or repulsion is essential for characteristics of the motion, namely, its stability properties. Motivation for these studies arises from the problem of motion and acceleration of matter (charged particles or dust grains)~\\cite{Sen:1997,Vok-Kar:1991a,Fel-Sor:2003,Cal-Fel-Fab-Tur:1982,Stu-Hle:1998}.  Compact stars can be imagined as being endowed with magnetic dipoles that are anchored in the stars and co-rotate with them~\\cite{Mestel:1999,Sen:1997,Rez-Ahm-Mil:2001}. Rotation gives rise to the electric component of the field, and we assume that  the resulting structure dominates over the small-scale turbulent fields. On the other hand, uncharged black holes do not support their own intrinsic magnetic fields; they can only be embedded in external fields~\\cite{Pet:1974,Wal:1974}. In general, the gravitational and electromagnetic fields interact with each other and the result is determined through the Einstein-Maxwell equations. To avoid such a complex description, one can also consider a test electromagnetic field influenced by the given gravitational field.  In both cases, magnetic compact stars as well as magnetized black holes\\footnote{We use the term magnetic compact stars (magnetic stars) to specify neutron stars, Q-stars and hybrid stars as well, and the magnetized black holes for the black holes embedded in ordered magnetic fields.}, the interplay of strong gravity, rotation and magnetic field is capable of accelerating electrically  charged particles and, possibly, it can also introduce new families of  stable motion. The latter will be the main subject of the present paper.  Recently, we have studied the existence of stable circular orbits with fixed latitudes of charged particles moving off the equatorial plane (halo orbits)~\\cite{Kov-Stu-Kar:2008,Stu-Kov-Kar:2009}. We assumed an axially symmetric (aligned)  magnetic dipolar field and built our discussion on classical studies~\\cite{How-Hor-Ste:1999,Dull-Hor-How:2002} derived in the context  of weak (Newtonian) gravitational fields, and designed for planetary studies. Particles on halo orbits are bound to the central object by a combined effect of gravitational and electromagnetic  forces and they do not have to cross neither the equatorial plane nor  the rotation axis. In the poloidal plane, the halo orbits are located in a kind of lobes that are positioned symmetrically above and below the equatorial  plane. The lobes may be disjoint or they may be interconnected,  depending on the parameters. Our study~\\cite{Kov-Stu-Kar:2008}, within  the pseudo-Newtonian approach~\\cite{Stu-Kov:2008,Pac-Wii:1980,Abr:2009}, has shown that the stable halo orbits  can indeed emerge also near magnetic compact stars, where their structure is affected by strong gravitational field. We have also examined the case of  charged and rotating (Kerr-Newman) black holes and naked singularities.  Surprisingly, the stable halo orbits do not exist above the horizon of Kerr-Newman black holes. This suggests that the structure of the Kerr-Newman spacetime is in several aspects quite special one, while the generic case of dipole-type field usually allows for the motion along stable halo orbits.  Here, we extend our investigation of the halo orbits to the case of magnetic compact stars and magnetized black holes in general relativity.  To this end, we study the existence and orientation of stable halo orbits which surround the minima of two-dimensional effective potential in the poloidal ($r,\\theta$)-plane. We also deal with qualitatively different cases of the general, in the halo lobes bounded, off-equatorial motion (halo motion). Namely, we point out to the motion in two separated halo lobes extending symmetrically above and below the equatorial plane versus the motion in a single region arising from the merged lobes across the equatorial plane. We also distinguish the cases when the halo lobes admit outflow of particles to the outer space or they allow the inflow onto the central body.  In the space environment, dust particles can carry small electrostatic charges~\\cite{Hor:1996}. An astrophysical motivation to study the halo motion of such weakly charged particles is the search for circumpulsar debris that could be formed from a fall-back disc resulting from the fraction of the explosion ejecta material that fails to escape~\\cite{Cordes08,Muno06,Wang06}. The anomalous X-ray pulsars~\\cite{dunc92} represent a category of magnetars around which the infrared observation indicate the presence of debris dust discs, potentially relevant for our deliberation. These are young neutron stars in the category of magnetars with slow spin periods in a range $\\approx 2$--$12$ seconds. The minimum distance to which the dust can reach is uncertain, however, part of it may enter the magnetosphere and influence the current flows below the light cylinder, until the dust evaporates. The potential relevance of the halo motion of dust particles is based on the fact that these orbits occupy areas of the stable motion. The irradiated dust particles acquire a small net electric charge and modulate the source radiation at their characteristic oscillation frequencies as they move in the halo lobes.  The paper is organized as follows. In section \\ref{Sec:Formalism}, we combine the standard formulation based  on the construction of the super-Hamiltonian and the corresponding effective potential, with perhaps a less frequent (in this context) formalism of forces in the projected three-space~\\cite{Abr-Nur-Wex:1995,Agu-etal:1996,Stu-Hle-Jur:2000,Kov-Stu:2007}. We maintain the assumption about axial symmetry but we generalise ref.~\\cite{Kov-Stu-Kar:2008} by considering the background of the magnetic compact star in a more consistent framework, i.e. in the Schwarzschild geometry with a test dipole-type rotating magnetic field~\\cite{Sen:1997} (section~\\ref{Sec:MagneticStar}). Furthermore, as a qualitatively different situation, we also consider a rotating black hole immersed in the uniform magnetic field~\\cite{Wal:1974} (section {\\ref{Sec:Wald}). We discuss the astrophysical relevance in section~\\ref{sec:Relevance} and then we conclude the paper in section~\\ref{sec:Conclusions}. ",
        "conclusions": "Conclusions} Our calculations reveal the properties of halo motion of charged particles near magnetic compact stars and black holes. These can be relevant e.g. for charged dust grains when they acquire a small magnitude of electric charge and occur in the strong gravitational field. Such particles can originate from supernova material which falls back and migrates into the pulsar vicinity before evaporating.  The halo orbits can take place in axially symmetric systems: near a magnetic compact star, which we modelled by the Schwarzschild geometry and a test rotating dipole magnetic field, as well as near a Kerr black hole immersed in an asymptotically uniform magnetic field.  In this paper, we concentrated ourselves on the methodological aspects of the halo orbits and their systematic classification. Our main interest concerned the properties of the halo orbits in situations when the motion takes place in a strong gravitational field. We only briefly touched the astrophysical applications and will investigate them further in the following work. Table \\ref{Tab:3} summarizes the basic characteristics of different  examples of strongly gravitating systems which exhibit the halo orbits. \\begin{table}[tbh!] \\caption{Orientation of motion of electrically charged particles along stable halo orbits and the corresponding sign of the charge. The results correspond to the configuration where the magnetic field (characterized by the magnetic  dipole moment $\\mathcal{M}$ or magnetic field strength in the equatorial  plane $B_0$) is oriented in the same direction as the rotation  vector, characterized by the signs of $\\Omega$ or $a$.}   {\\small\\vspace*{4mm} \\begin{indented} \\item\\begin{tabular}{@{}lll} \\br\\rule[-1em]{0pt}{2.5em} Compact object&Charge of particles&Orientation of motion\\\\ \\hline \\hline\\rule[0em]{0pt}{1.5em} Rotating magnetic star & Positive & Co-rotation\\\\ & Positive & Counter-rotation\\\\ \\rule[-1em]{0pt}{1.5em}                     & Negative  & Co-rotation\\\\ \\hline\\rule[0em]{0pt}{1.5em} Static magnetic star & Positive & Counter-rotation\\\\ \\rule[-1em]{0pt}{1.5em}   & Negative  & Co-rotation\\\\ \\hline\\rule[0em]{0pt}{1.5em} Magnetized black hole & Positive & Counter-rotation\\\\ \\rule[-1em]{0pt}{1.5em} (Kerr metric) & Negative  & No-orbits\\\\ \\br \\end{tabular} \\end{indented} \\label{Tab:3}} \\end{table}  As for the magnetic compact stars, results of our discussion of the stable halo orbits orientation are qualitatively consistent with those obtained previously in Newtonian and the pseudo-Newtonian investigations~\\cite{Dull-Hor-How:2002,Kov-Stu-Kar:2008}. We notice that the strong gravitational field and the presence of event horizon affect the form of the effective potential and introduce significant changes  with respect to the Newtonian case. However, we also notice that the  pseudo-Newtonian approach can give surprisingly good information on the loci of halo orbits in spherically symmetric gravitational fields of magnetic stars, in which case the results qualitatively agree with the  exact general relativistic approach presented here. On the other hand, the pseudo-Newtonian modelling of Kerr spacetimes~\\cite{SemK99,Muk:2002}  is much more complicated, and seems to be less fruitful for modelling the phenomena in the Kerr geometry, which itself is relatively simple.  The effective potential of the halo motion exhibits a number of qualitatively different types. Two basic categories can be distinguished. The first one allows particles, bounded within the halo potential lobes, to commute across the equatorial plane. On the other hand, the second category does not allow such an interconnection, as the two lobes are entirely disjoint. The off-equatorial halo motion is an interesting phenomenon not only because it has not yet been explored in full detail in the literature, but also because it shows a variety of orbits, depending on orbital parameters: off-equatorial stable circular motion which does not cross the equatorial plane; small oscillations in the radial and vertical directions, with characteristic frequencies around the circular orbits; and the motion that traverses between the lobes when the energy is increased to sufficiently high levels. Recently, the off-equatorial motion of charged particles near Kerr black hole in asymptoticaly uniform magnetic field has been discussed from a different point of view in the ref. \\cite{Pre:2010} as well.   \\ack The authors thank M. Urbanec for useful discussions concerning section 5.1. Institute of Physics and the Astronomical Institute have been operated under the projects  MSM\\,4781305903 and AV\\,0Z10030501, and further supported by  the Centre for Theoretical Astrophysics LC06014 in the Czech Republic. JK, VK and ZS thank the Czech Science Foundation (ref.\\ P209/10/P190, 205/07/0052, 202/09/0772). OK acknowledges the doctoral student program of the Czech Science Foundation (ref.\\ 205/09/H033)."
    },
    "1005/1005.3793_arXiv.txt": {
        "abstract": "AR Aur is the only eclipsing binary known to contain a HgMn star, making it an ideal case for a detailed study of the HgMn phenomenon. HgMn stars are a poorly understood class of chemically peculiar stars, which have traditionally been thought not to possess significant magnetic fields. However, the recent discovery of line profile variability in some HgMn stars, apparently attributable to surface abundance patches, has brought this belief into question. In this paper we investigate the chemical abundances, line profile variability, and magnetic field of the primary and secondary of the AR Aur system, using a series of high resolution spectropolarimetric observations. We find the primary is indeed a HgMn star, and present the most precise abundances yet determined for this star. We find the secondary is a weak Am star, and is possibly still on the pre-main sequence. Line profile variability was observed in  a range of lines in the primary, and is attributed to inhomogeneous surface distributions of some elements. No magnetic field was detected in any observation of either stars, with an upper limit on the longitudinal magnetic field in both stars of 100~G. Modeling of the phase-resolve longitudinal field measurements leads to a $3\\sigma$ upper limit on any dipole surface magnetic field of about 400~G. ",
        "introduction": "HgMn stars are chemically peculiar late B stars, characterized by strong overabundances of Mn, by up to a thousand times solar, and strong overabundances of Hg, by up to a hundred thousand times solar. Strong overabundances of Xe, iron-peak elements, Ga, and underabundances of He are also often seen in these stars. These chemical anomalies are believed to be produced by the selective radiative levitation and gravitational settling of chemical elements in the outer stellar layers \\citep*[e.g.][]{Michaud1974}. However, despite a significant number of observational and theoretical studies, the physical processes responsible for % these strong chemical peculiarities remain poorly understood. HgMn stars have been generally thought not to posses strong (dynamically important) magnetic fields \\citep[e.g.][]{Shorlin2002}, setting them apart from magnetic chemically peculiar Ap and Bp stars.  Recently, line profile variability has been detected in some spectral lines of a number of HgMn stars. This has been interpreted as an inhomogeneous surface distribution of specific elements (\\citealt*{Ryabchikova1999-alphaAnd-abun-var}; \\citealt{Adelman2002-HgMn-var-1st,Kochukhov2005-hgmn-inhomo}). While magnetic chemically peculiar Ap and Bp stars display a wide range of surface abundance inhomogeneities, HgMn stars are the only type of (apparently) non-magnetic A or B stars in which such inhomogeneities have been found.  In Ap and Bp stars, surface abundance inhomogeneities are usually attributed to the magnetic field, possibly working together with stellar rotation and magnetically-channelled mass loss. This has led some authors to suggest that HgMn stars may have hitherto undetected magnetic fields, which could give rise to the observed abundance inhomogeneities \\citep[e.g.][]{Adelman2002-HgMn-var-1st,Hubrig2006-arAur-lVar}. A detailed examination of the spotted HgMn star $\\alpha$~And by \\citet{Wade2006-alphaAnd} found no evidence for such a magnetic field, with an upper limit on the dipole field strength of between 50 and 100 G. Notwithstanding the apparent absence of magnetic fields in $\\alpha$~And, the more general presence of magnetic fields in variable HgMn stars has not been strongly tested.  Doppler imaging of surface abundance inhomogeneities has been performed by \\citet{Adelman2002-HgMn-var-1st} and \\citet{Kochukhov2007-alphaAnd-map-var} for the HgMn star $\\alpha$~And. Surprisingly \\citet{Kochukhov2007-alphaAnd-map-var} find secular evolution of the Doppler maps on a time scale of years. This is the only known case of surface abundance inhomogeneities evolving with time in any A or B star.  AR Aurigae (17 Aur, HR\\,1728, HD\\,34364) is a unique double-lined spectroscopic binary system with nearly identical components (B9V+B9.5V). Being the only eclipsing binary known to host a HgMn star, it allows for a precise, model independent determination of fundamental physical parameters. The primary of the system has been long known as an HgMn star \\citep{Wolff1978} and line variability in the primary has recently been reported by \\citet{Hubrig2006-arAur-lVar}. This makes the system an ideal candidate for a more detailed investigation.  The two main components of the AR Aur system have similar masses of 2.48~\\Msun\\, and 2.29~\\Msun, and an orbital period of 4.135 days \\citep{Nordstrom1994-arAur-fundamental}. \\citet{Chochol1988-araur-triple} discovered a third component to the AR Aur system, based on the light-time effect. According to \\citet*{Albayrak2003}, the third star has a mass of 0.54~\\Msun, a separation of 13~AU, and a period of 23.7 years. \\citet{Nordstrom1994-arAur-fundamental} performed a detailed study of the photometric light curves, deriving radii, masses, and effective temperatures of the two brighter, eclipsing components of AR Aur, as well as considering the age of the system. They suggest that the secondary is likely a pre-main sequence star still contracting towards the zero-age main sequence line (ZAMS), while the primary is at the ZAMS.  \\citet{Ryabchikova1998-HgMn-binaries} performed an abundance analysis of both components of AR Aur, based on the equivalent widths from photographic spectra reported by \\citet{Khokhlova1995-araur-ew}, and found strong overabundances of a wide range of elements in the primary, including Mn, Fe, Sr, Y, Pt, and Hg, typical of HgMn stars. Inconclusive evidence of the line profile variability in AR Aur A was given by \\citet*{Takeda1979} and \\citet*{Zverko1997} based on low-quality photographic spectra. \\citet{Hubrig2006-arAur-lVar} studied AR Aur with modern spectroscopic instrumentation, demonstrating line variability for several elements and suggesting inhomogeneous surface distributions for Sr and Y.  In this work we search for magnetic fields in AR Aur A and B, making this system the second HgMn star with very precise upper limits on its magnetic field. Additionally, we derive improved chemical abundances for the primary and the secondary, confirming the primary as a prominent HgMn star, and discovering that the secondary is a weak Am star. ",
        "conclusions": "We find AR Aur A to be a strongly peculiar HgMn star, confirming previous results \\citep{Khokhlova1995-araur-ew,Ryabchikova1998-HgMn-binaries}. AR Aur B shows weak Am peculiarities, particularly modest overabundances of Fe-peak elements and Nd, and modest underabundances of Ca and Sc.  Thus AR Aur is a binary system with two different chemically peculiar stars. This is striking in light of the very similar temperatures of the two stars. It is possible that the small temperature difference represents a sharp dividing line between Am and HgMn stars. Alternately, if AR Aur B is still on the pre-main sequence, the difference in chemical abundances may be a result of the different evolutionary status of the stars.  We confirm the result of \\citet{Nordstrom1994-arAur-fundamental} that the cooler, lower mass, secondary has a larger radius than the primary, assuming that both components have the same Mg abundance. Under this assumption, we agree with their conclusion that this implies the secondary is still contracting towards the ZAMS, while the primary has likely already arrived on the ZAMS.  Thus both components of AR Aur are likely some of the youngest chemically peculiar stars known. AR Aur B may be the first pre-main sequence Am star discovered. This result, if correct, implies that chemical peculiarities must arise quickly in both HgMn stars and in Am stars.  The atomic diffusion process thought to drive the observed chemical peculiarities must be fairly efficient for this to occur.  Despite a careful high S/N search, we find no evidence for a magnetic field in AR Aur A or B. The LSD profiles for our observations, which are sensitive to simpler field geometries with a zero net longitudinal component, as well as more complex field geometries, show no indication of any signal in either star. Our longitudinal field measurements place a $3\\sigma$ upper limit of 100~G on the longitudinal fields of both stars over a range of rotational phases. This constrains a dipole field to be less than about 400 G at the pole in either star. Thus we conclude that neither AR Aur A nor B is likely to have any significant magnetic field. This matches the result of \\citet{Wade2006-alphaAnd} for the HgMn star $\\alpha$~And.  AR Aur represents a second case in which line profile variability and surface abundance inhomogeneities exist in the absence of any strong magnetic field. This substantially strengthens the conclusion that, whatever the mechanism producing these inhomogeneities, a magnetic field is not required."
    },
    "1005/1005.1105_arXiv.txt": {
        "abstract": "s{ The statistics of peaks of the initial, Gaussian density field  can be used to interpret the abundance and clustering of massive dark matter haloes. I discuss some recent theoretical results related to their clustering and its redshift evolution. Predictions from the peak model are qualitatively consistent with measurements of the linear bias  of high mass haloes, which also show some evidence for a dependence on the halo mass $M$ at fixed peak height $\\nu$. The peak approach also predicts a distinctive scale-dependence in the bias of haloes across the baryon acoustic feature, a measurement of which would provide strong support for its validity. For 2$\\sigma$ density peaks collapsing at $z=0.3$, this residual scale-dependent bias is at the 5-10\\% level and should thus be within reach of very large simulations of structure formation. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "1005/1005.2541.txt": {
        "abstract": "{The Low Frequency Instrument (LFI) on-board the ESA Planck satellite carries eleven radiometer subsystems, called Radiometer Chain Assemblies (RCAs), each composed of a pair of pseudo-correlation receivers. We describe the on-ground calibration campaign performed to qualify the flight model RCAs and to measure their pre-launch performances. Each RCA was calibrated in a dedicated flight-like cryogenic environment with the radiometer front-end cooled to 20K and the back-end at 300K, and with an external input load cooled to 4K. A matched load simulating a blackbody at different temperatures was placed in front of the sky horn to derive basic radiometer properties such as noise temperature, gain, and noise performance, e.g.  1/f noise. The spectral response of each detector was measured as was their susceptibility to thermal variation. All eleven LFI RCAs were calibrated. Instrumental parameters measured in these tests, such as noise temperature, bandwidth, radiometer isolation, and linearity, provide essential inputs to the Planck-LFI data analysis.} ",
        "introduction": "The Planck mission\\footnote{Planck \\emph{(http://www.esa.int/Planck)} is a project of the European Space Agency - ESA - with instruments provided by two scientific Consortia funded by ESA member states (in particular the lead countries: France and Italy) with contributions from NASA (USA), and telescope reflectors provided in a collaboration between ESA and a scientific Consortium led and funded by Denmark.} has been developed to provide a deep, full-sky image of the cosmic microwave background (CMB) in both temperature and polarization. Planck incorporates an unprecedented combination of sensitivity, angular resolution and spectral range - spanning from centimeter to sub-millimeter wavelengths - by integrating two complementary cryogenic instruments in the focal plane of the Planck telescope. The Low Frequency Instrument (LFI) covers the region below the CMB blackbody peak in three frequency bands centered at 30, 44 and 70 GHz. The spectral range of the LFI is also suitable for a wealth of galactic and extragalactic astrophysics. The LFI maps will address studies of diffuse Galactic free-free and synchrotron emission, emission from spinning dust grains, and  discrete Galactic radio sources. Extragalactic radio sources will also be observed, particularly those with flat or strongly inverted spectra, peaking at mm wavelengths. Furthermore, the Planck scanning strategy will also allow monitoring of radio source variability on a variety of time scales. While these are highly interesting astrophysical objectives, the LFI design and calibration is driven by the main Planck scientific scope, i.e., CMB science.  Low Frequency Instrument is an array of cryogenic radiometers based on indium phospide (InP) cryogenic HEMT low noise amplifiers \\citep{bersanelli2009}. The array is composed of 22 pseudo--correlation radiometers mounted in eleven independent radiometer units called \"radiometer chain assemblies\" (RCAs), two centered at 30 GHz, three at 44 GHz and six at 70 GHz \\footnote{The chains are numbered as {\\tt RCAXX} where {\\tt XX} is a number from {\\tt 18} to {\\tt 23} for the 70 GHz RCAs, from {\\tt 24} to {\\tt 26} for the RCAs at 44 GHz, and from {\\tt 27} to {\\tt 28} for the RCAs at 30 GHz}. To optimize the sensitivity and minimize the power dissipation in the front end, each RCA is split into a front-end module (FEM), cooled to 20K, and a back-end module (BEM), operating at 300K, connected by a set of composite waveguides.  Accurate calibration is mandatory for optimal operation of the instrument during the full-sky survey and for measuring parameters that are essential for the Planck data analysis.  The LFI calibration strategy has been  based on a complementary approach that includes both pre-launch and post-launch activities. On-ground measurements were performed at all-unit and sub-unit levels, both for qualification and performance verification. Each single FEM and BEM as well as the passive components (feed horn, orthomode transducers, waveguides, 4K reference loads) were tested in a stand--alone configuration before they were integrated into the RCA units.  The final scientific calibration of the LFI was carried out at two different integration levels, depending on the measured parameter. First each RCA was tested independently in dedicated cryofacilities, which are capable of reaching a temperature of $\\sim$4K at the external input loads. These conditions were necessary for an accurate measurement of key parameters such as system noise temperature, bandwidth, radiometer isolation, and linearity. Subsequently, the eleven RCAs were integrated into the full LFI instrument (the so-called \u00d2radiometer array assembly\u00d3, RAA) and tested as a complete instrument system in a large cryofacility, with highly stable input loads cooled down to 20K.  We reports on the calibration campaign of the RCAs, while \\citet{mennella2009a} report on the RAA calibration. The parameters derived here are crucial for the Planck-LFI scientific analysis. Noise temperatures, bandwidths, and radiometer isolation provide essential information to construct an adequate noise model, which is needed as an input to the map-making process. Any non-linearity of the instrument response must be accurately measured, because corrections may be needed in the data analysis, particularly for observations of strong sources such as planets (crucial for in-flight beam reconstruction) and the Galactic plane. As part of the RCA testing, we also performed an end-to-end measurement of the bandshape inside the cryofacility. Finally, for comparison and as a consistency check, the RCA test plan also included measurement of parameters whose primary calibration relies on the RAA test campaign, such as optimal radiometer bias (tuning), 1/f noise (knee frequency and slope), gain, and thermal susceptibility.  The 30 and 44 GHz RCAs were integrated and tested in Thales Alenia Space Italia (TAS-I), formerly Laben, from the beginning of January 2006 to end of May 2006 using a dedicated cryofacility \\citep{terenzi2009a} to reproduce flight-like thermal interfaces and input loads. The 70 GHz RCAs were integrated and calibrated in Yilinen Electronics (Finland) from the end of April 2005 to mid February 2006 with a similar cryofacility, but with simplified thermal interfaces \\citep{terenzi2009a}. The differences between the two cryofacilities resulted in slightly different test procedures, because the temperatures were not controlled in the same way.  Section \\ref{Concept} of the paper describes the concept of RCA calibration, while Sect. \\ref{cryo} illustrates the two cryofacilities. In Sect. \\ref{results} we describe each test, the methods used in the analysis and the results. The conclusions are given in Sect. \\ref{conclusions}. ",
        "conclusions": "The eleven LFI Radiometer Chain Assemblies  were calibrated according to the overall LFI calibration plan. The front end low-noise amplifiers and phase switches were properly tuned to guarantee minimum noise temperature and best isolation.  Basic performances (noise temperature, isolation, linearity, photometric gain), noise properties (1/f spectrum, noise equivalent bandwidth), relative bandshape, and susceptibility to thermal variations were measured in a dedicated cryogenic environment as close as possible to flight-operational conditions. All radiometric parameters were measured with excellent repeatability and reliability, except for 1/f noise at 70 GHz and some of the bandpasses at 30 GHz. The measurements were essentially in line with the design expectations, indicating a satisfactory instrument performance. Ultimately all RCA units were accepted, because the measured performances were in line with the scientific goal of LFI. The RCA test campaign described here represents the primary calibration test for some key radiometric parameters of the LFI, because they were not accurately measured as part of the RAA calibration campaign, nor can they be measured in-flight during the Planck nominal survey: \\begin{itemize} \\item tuning results were used to set the subsequent tuning verification procedure up to the calibration performance and verification phase (CPV) in flight; \\item non-linear behavior of the 30 and 44 GHz RCAs was accurately measured and characterized and used to estimate the impact in flight \\citep{mennella2009b}. Moreover each radiometer system noise temperature was accurately determined; \\item except for some of the measured spectral responses, in which systematic effects arising from standing waves in the sky load were experienced,  the RCA band shapes were only measured and characterized during the RCA tests. Together with the independent estimates of the band shape based on a synthesis of measured responses at unit level \\citep{zonca2009},  they are essential for the flight data analysis; \\item susceptivity  to thermal variation of the FEM and BEM was measured and represents the reference values because only the FEM susceptivity was measured during RAA tests \\citep{terenzi2009b}, but under worse thermal conditions. \\end{itemize}  In conclusion we can state that even if the RCA calibration campaign was an intermediate step in the LFI development, the results obtained and presented here will be used in conjunction with the performance measured in flight to the exploitation of the scientific goal of Planck."
    },
    "1005/1005.0386_arXiv.txt": {
        "abstract": "We explore a technique for identifying the highest redshift ($z>4$) sources in {\\it Herschel}/SPIRE and BLAST submillimeter surveys by localizing the position of the far-infrared dust peak. Just as {\\it Spitzer}/IRAC was used to identify stellar `bump' sources, the far-IR peak is also a redshift indicator; although, the latter also depends on the average dust temperature. We demonstrate the wide range of allowable redshifts for a reasonable range of dust temperatures and show that it is impossible to constraint the redshift of individual objects using solely the position of the far-IR peak. By fitting spectral energy distribution models to simulated {\\it Herschel}/SPIRE photometry we show the utility of radio and/or far-infrared data in breaking this degeneracy. With prior knowledge of the dust temperature distribution it is possible to obtain statistical samples of high redshift submillimeter galaxy candidates. We apply this technique to the BLAST survey of ECDFS to constrain the number of dusty galaxies at $z>4$. We find $8\\pm2$ galaxies with flux density ratios of $S_{500}>S_{350}$; this sets an upper limit of $17\\pm4\\,$deg$^{-2}$ if we assume all are at $z>4$. This is $<35\\,$\\% of all 500$\\,\\mu$m-selected galaxies down to S$_{500}>45\\,$mJy ($L_{\\rm{IR}}>2\\times10^{13}L_{\\odot}$ for $z>4$). Modeling with conventional temperature and redshift distributions estimates the percentage of these 500$\\,\\mu$m peak galaxies at $z>4$ to be between 10--85\\%. Our results are consistent with other estimates of the number density of very high redshift submillimeter galaxies and follows the decline in the star formation rate density at $z>4$. ",
        "introduction": "\\label{sec:intro}  The prevalence of dusty, infrared luminous galaxies at early times tests our theories about early galaxy formation and the production of metals. Semi-analytic models predict very small numbers of submillimeter galaxies (SMGs) at extreme redshifts ($z>4$, e.g.~Baugh et al.~2005), and the best observational constraints on the redshift distribution of SMGs are consistent with this (Chapman et al.~2005; Swinbank et al.~2009). Much attention has been given to recent studies which have spectroscopically confirmed a small number of SMGs at $z>4$ (Capak et al.~2008; Coppin et al.~2009; Daddi et al.~2009ab; Knudsen et al.~2010); however, these sources are still only a small fraction of the submillimeter population. Small number statistics and cosmic variance in current submillimeter surveys prohibit an accurate measure of the space density of the most distant SMGs.  The question of the prominence of the high redshift tail of dusty SMGs can be addressed with imminent submillimeter wide area surveys such as those with the {\\it Herschel Space Observatory} (Pilbratt 2001), the SCUBA-2 camera on the JCMT (Holland et al.~2006), and the AzTEC camera on the Large Millimeter Telescope (LMT, Wilson et al.~2008). The {\\it Herschel Space Observatory} is currently collecting data from 70--500$\\,\\mu$m over wide survey fields. With an aperture of 3.5$\\,$m, {\\it Herschel} is only sensitive to ultra-luminous infrared galaxies (ULIRGs, L$_{IR}>10^{12}L_{\\odot}$) above $z>2$ at 250--500$\\,\\mu$m before it reaches the confusion limit. One advantage of SPIRE surveys is the simultaneous observations at three wavelengths; 250, 350 and 500$\\,\\mu$m, which samples the peak of the far-IR spectral energy distribution (SED) due to dust emission in high redshift galaxies. With prior knowledge of the dust properties which describe the shape of the SED, the position of the far-IR peak is sensitive to the redshift of the galaxy.  In this paper, we discuss the power and limitations of using the far-IR peak as a redshift indicator in submillimeter surveys. We apply this technique to the current BLAST survey of ECDFS and place an upper limit on the number density of 500$\\,\\mu$m-selected SMGs at $z>4$.  Throughout this paper we assume a standard cosmology with $H_{0}=71\\,\\rm{km}\\,\\rm{s}^{-1}\\,\\rm{Mpc}^{-1}$, $\\Omega_{\\rm{M}}=0.27$ and $\\Omega_{\\Lambda}=0.73$. ",
        "conclusions": "From the BLAST ECDFS data we estimate a space density of $<17\\,$deg$^{-2}$ for sources with S$_{500}>45\\,$mJy at $z>4$; this is $<35\\,$\\% (8/23) of all sources selected down to these depths in the BLAST data. These bright 500$\\,\\mu$m-selected SMGs at $z>4$ contribute $<1\\%$ of the total background emission at 500$\\,\\mu$m (Fixsen et al.~1998).  Constrained by a recent analysis of the infrared extragalactic background light (EBL), galaxies evolution models predict the source density of galaxies with S$_{500}>27\\,$mJy\\footnote{In order to compare to the models we need to assume a rough correction for flux boosting of the BLAST measurements. Based on Eales et al.~(2009), we adopt an average correction of 0.6 to the BLAST fluxes which effectively brings the flux limit down to 27$\\,$mJy.} at $z>4$ to be 1--5$\\,$deg$^{-2}$ (Chary \\& Pope 2010). The range of predicted number density comes from considering two evolutions of the models both of which fit the observed EBL. We find that our current observational constraint is consistent with both evolutionary scenarios. Future submm surveys with much larger telescopes such as JCMT, LMT and CCAT will reach much deeper flux limits before hitting the confusion limit. With these surveys, we can expect to detect many more dusty galaxies at $z>4$; the models in Chary \\& Pope (2010) predict a number density of 40-200$\\,$deg$^{-2}$ at $z>4$ for galaxies down to 5$\\,$mJy at 500$\\,\\mu$m.  In the past two years, the tally of $z>4$ spectroscopically confirmed SMGs has grown from zero to five\\footnote{A sixth $z>4$ SMG recently appeared in the literature (Knudsen et al.~2010).} (Capak et al.~2008; Coppin et al.~2009; Daddi et al.~2009a,b). Coppin et al.~(2009) estimate the number density of SMGs at $z>4$ by combining these 5 sources from 3 independent submillimeter surveys and find a lower limit of $>7\\,$deg$^{-2}$. While this is consistent with our predictions for 500$\\,\\mu$m-selected galaxies we note that the sensitivities of the BLAST and submillimeter surveys are not the same at $z=4$; only $1/5$ of the $z>4$ spectroscopically confirmed SMGs might be detected down to the BLAST ECDFS survey depths (assuming a typical SMG dust temperature). Folding this in, the Coppin et al.~(2009) estimate down to S$_{850}>20\\,$mJy scales to $>1.4\\,$deg$^{-2}$ at $z>4$.   There is some evidence that large numbers of very massive galaxies at $z>4$ pose a substantial challenge to standard galaxy formation theory (e.g.~Baugh et al.~2005). Fig.~\\ref{fig:num} summarizes the current observational constraints on the number density of bright SMGs as a function of redshift. Blue and green bars are for 850$\\,\\mu$m-selected SMGs using values from the literature (Chapman et al.~2005; Coppin et al.~2009; Pope et al.~2006) and the red bar is our new upper limit from the analysis in this paper of the BLAST 500$\\,\\mu$m-selected sources. Overall, there does not appear to be room for a large fraction of the SMG population to be above $z=4$ down to these flux limits."
    },
    "1005/1005.1455_arXiv.txt": {
        "abstract": "The Palomar Transient Factory (PTF) is systematically charting the optical transient and variable sky. A primary science driver of PTF is building a complete inventory of transients in the local Universe (distance less than 200 Mpc). Here, we report the discovery of PTF\\,10fqs, a transient in the luminosity ``gap'' between novae and supernovae. Located on a spiral arm of Messier~99, PTF\\,10fqs has a peak luminosity of $M_r = -12.3$, red color ($g-r=1.0$) and is slowly evolving (decayed by 1\\,mag in 68\\,days). It has a spectrum dominated by intermediate-width H$\\alpha$ ($\\approx$930\\,km\\,s$^{-1}$) and narrow calcium emission lines. The explosion signature (the light curve and spectra) is overall similar to that of M85\\,OT2006-1, SN\\,2008S, and NGC\\,300\\,OT.  The origin of these events is shrouded in mystery and controversy (and in some cases, in dust). PTF\\,10fqs shows some evidence of a broad feature (around 8600\\,\\AA) that may suggest very large velocities ($\\approx$10,000\\,km\\,s$^{-1}$) in this explosion. Ongoing surveys can be expected to find a few such events per year. Sensitive spectroscopy, infrared monitoring and statistics (e.g. disk versus bulge) will eventually make it possible for astronomers to unravel the nature of these mysterious explosions. ",
        "introduction": "Two reasons motivate us to search for transients in the local Universe (distance $<$ 200\\,Mpc). First, the emerging areas of gravitational wave astronomy, high-energy cosmic rays, very high-energy photons, and neutrino astronomy are limited to this distance horizon either due to physical effects (optical depth) or instrumental sensitivity. Thus, to effectively search for an electromagnetic analog, understanding the full range of transient phenomena is essential. For instance, the electromagnetic counterpart to the gravitational wave signature of neutron star mergers is expected to be fainter and faster than that of supernovae (e.g. \\citealt{mmd+10}).  Our second motivation is one of pure exploration.  The peak luminosity of novae ranges between $-4$ and $-10$\\,mag,\\footnote{Unless explicitly noted, quoted magnitudes are in the $R$ band} whereas supernovae range between $-15$ and $-22$\\,mag. The large gap between the cataclysmic novae and the catastrophic supernovae has been noted by early observers. Theorists have proposed several intriguing scenarios producing transients in this ``gap.'' (e.g. \\citealt{bsw+07,mpq+09,skw+10,mtt+10})  The Palomar Transient Factory\\footnote{\\texttt{http://www.astro.caltech.edu/ptf.}} (PTF; see \\citealt{gsv+08,lkd+09,rkl+09}) was designed to undertake a systematic exploration of the transient sky in the optical bands. One of the key projects of PTF is to build a complete inventory of transients in the local Universe. PTF has a ``Dynamic'' cadence experiment which undertakes frequent observations of fields, optimized for inclusion of galaxies in the local Universe.  A description of the design sensitivity is given elsewhere \\citep{kk09}.  Here, we report on the discovery of PTF\\,10fqs, a transient in this ``gap'' between novae and supernovae. ",
        "conclusions": "PTF\\,10fqs is the fourth member of a class of extragalactic transients\\footnote{Henceforth we use the term ``luminous red novae'' as a functional short name for such events.} which possess a peak luminosity between that of novae and supernovae, and have spectral and photometric evolution that bear no resemblance to either supernovae or novae. The other members of this class are M85-OT, NGC\\,300-OT, SN\\,2008S.  NGC\\,300-OT and SN\\,2008S are remarkable for their very bright mid-infrared progenitors. Though sensitive pre-explosion observations of M85-OT and PTF\\,10fqs do exist, the large distance to the Virgo Cluster (17\\,Mpc) relative to that of NGC\\,300 (1.9\\,Mpc) and NGC6946 (5.7\\,Mpc) results in weak constraints on the luminosity of any pre-explosion star. PTF\\,10fqs, NGC\\,300-OT, and SN\\,2008S occurred in star-forming regions whereas M85-OT was in the bulge. {\\it Prima facie}, this group of explosive events can be divided into a disk and a bulge group.  The discovery of PTF\\,10fqs in itself cannot address whether the two groups of luminous, red novae are one and the same. The proposed models to explain this group are diverse: electron capture within an extreme asymptotic giant branch (AGB) star, common-envelope phase (stellar merger), inspiral of a giant planet into the envelope of an aging parent star, a most peculiar nova, and a most peculiar supernova.  The possible evidence of the broad feature centered around the Ca~II near-IR triplet with an inferred velocity spread of 10,000\\,km s$^{-1}$ may be an important clue. It would mean that these events possess both a low- and a high-velocity outflow. By comparison with other astronomical sources, one can envisage a high-velocity polar outflow and a slower equatorial outflow (but with a larger mass).  To this end, continued sensitive spectroscopy of PTF\\,10fqs (and of course other such future events) would be very valuable.  The ``Transients in the Local Universe'' key project of the Palomar Transient Factory is designed to systematically unveil events in the gap between novae and supernovae. It surveys $\\approx$20,000 nearby galaxies ($d < 200$\\,Mpc) yearly at 1-day cadence and a depth of $R < 21$ mag.  (If the maximum luminosity of this class is $-14$ mag, then we would be sensitive to events out to 100\\,Mpc.) Furthermore, {\\it Spitzer} has a growing archive of deep images of nearby galaxies (e.g., SINGS, \\citealt{kab+03}; LVL, \\citealt{dcj+09}, and S4G, \\citealt{srh+10}), and {\\it WISE} \\citep{wem+10} has an ongoing all-sky survey in the mid-IR. This will allow us to probe deeper in search of the pre-explosion counterpart and possibly present a new channel for discovery of luminous red novae.  The discovery of PTF\\,10fqs is only the harbinger of the uncovering of a large sample of such transients to unveil the nature of this new class of explosions.    \\bigskip \\smallskip  \\noindent{\\bf Acknowledgments.}  M.M.K. thanks the Gordon and Betty Moore Foundation for a Hale Fellowship in support of graduate study. The Weizmann Institute PTF participation is supported in part by the Israel Science Foundation via grants to AGY. The Weizmann-Caltech collaborative PTF effort is supported by the US-Israel Binational Science Foundation. AGY and MS are jointly supported by the ``making connections'' Weizmann-UK program. AGY further acknowledges support by a Marie Curie IRG fellowship and the Peter and Patricia Gruber Award, as well as funding by the Benoziyo Center for Astrophysics and the Yeda-Sela center at the Weizmann Institute. A.V.F.'s group and KAIT are supported by National Science Foundation (NSF) grant AST-0908886, the Sylvia \\& Jim Katzman Foundation, the Richard \\& Rhoda Goldman Fund, Gary and Cynthia Bengier, and the TABASGO Foundation; additional funding was provided by NASA through {\\it Spitzer} grant 1322321, as well as {\\it HST} grant AR-11248 from the Space Telescope Science Institute, which is operated by Associated Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. J.S.B. and his group are partially funded by a DOE SciDAC grant. E.O.O. and D.P. are supported by the Einstein fellowship. L.B. is supported by the National Science Foundation under grants PHY 05-51164 and AST 07-07633.   We are grateful to the staff of the Gemini Observatory for their promptness and high efficiency in attending to our TOO request. Likewise, we thank the staff of the Very Large Array and the Hobby Eberly Telescope. We acknowledge the following internet repositories: SEDS (Messier Objects) and GOLDMine (Virgo Cluster), Finally, as always, we are grateful to the librarians who maintain the ADS, the NED, and SIMBAD data systems.  The Hobby-Eberly Telescope (HET) is a joint project of the University of Texas at Austin, the Pennsylvania State University, Stanford University, Ludwig-Maximillians-Universit\\\"at M\\\"unchen, and Georg-August-Universit\\\"at G\\\"ottingen. The HET is named in honor of its principal benefactors, William P. Hobby and Robert E. Eberly. The Marcario Low-Resolution Spectrograph is named for Mike Marcario of High Lonesome Optics, who fabricated several optics for the instrument but died before its completion; it is a joint project of the Hobby-Eberly Telescope partnership and the Instituto de Astronom\\'{\\i}a de la Universidad Nacional Aut\\'onoma de M\\'exico. {\\it GALEX} (Galaxy Evolution Explorer) is a NASA Small Explorer, launched in 2003 April. We gratefully acknowledge NASA's support for construction, operation, and science analysis for the {\\it GALEX} mission, developed in cooperation with the Centre National d'Etudes Spatiales of France and the Korean Ministry of Science and Technology. PAIRITEL is operated by the Smithsonian Astrophysical Observatory (SAO) and was made possible by a grant from the Harvard University Milton Fund, the camera loan from the University of Virginia, and the continued support of the SAO and UC Berkeley. The Expanded Very Large Array is operated by the National Radio Astronomy Observatory, a facility of the NSF operated under cooperative agreement by Associated Universities, Inc.   \\bigskip"
    },
    "1005/1005.3461.txt": {
        "abstract": "% In 2007, a companion with planetary mass was found around the pulsating subdwarf B star V391~Pegasi with the timing method, indicating that a previously undiscovered population of substellar companions to apparently single subdwarf B stars might exist. Following this serendipitous discovery, the \\mbox{EXOTIME}\\footnote{\\texttt{http://www.na.astro.it/$\\sim$silvotti/exotime/}} monitoring program has been set up to follow the pulsations of a number of selected rapidly pulsating subdwarf B stars on time-scales of several years with two immediate observational goals:  1) determine $\\dot{P}$ of the pulsational periods $P$  2) search for signatures of substellar companions in O$-$C residuals due to periodic light travel time variations, which would be tracking the central star's companion-induced wobble around the center of mass.  These sets of data should therefore at the same time: on the one hand be useful to provide extra constraints for classical asteroseismological exercises from the $\\dot{P}$ (comparison with ''local'' evolutionary models), and on the other hand allow to investigate the preceding evolution of a target in terms of possible ''binary'' evolution by extending the otherwise unsuccessful search for companions to potentially very low masses.  While timing pulsations may be an observationally expensive method to search for companions, it samples a different range of orbital parameters, inaccessible through orbital photometric effects or the radial velocity method: the latter favours massive close-in companions, whereas the timing method becomes increasingly more sensitive towards wider separations.  In this paper we report on the status of the on-going observations and coherence analysis for two of the currently five targets, revealing very well-behaved pulsational characteristics in \\mbox{HS\\,0444+0458}, while showing \\mbox{HS\\,0702+6043} to be more complex than previously thought. % ",
        "introduction": "\\label{s:intro} % While $\\dot{P}$ measurements exist for a number of pulsating white dwarfs \\citep[e.g.][]{1991ApJ...378L..45K,2000ApJ...534L.185K,2005ApJ...634.1311K,1999ApJ...522..973C,2008A&A...489.1225C,2008ApJ...676..573M}, similar measurements have only been published for one pulsating subdwarf B star so far, \\object{V391~Pegasi} (hereafter \\object{HS\\,2201+2610}).  \\object{HS\\,2201+2610} was first discovered to be a rapidly pulsating subdwarf B star by \\citet{2001A&A...368..175O}. \\citet{2002A&A...389..180S,2007Natur.449..189S} were able to derive $\\dot{P}$ values for the two strongest pulsation modes, and found an additional pattern in the O$-$C diagrams that revealed the presence of a giant planet in a 3.2 year orbit. Additional slow pulsations were subsequently reported by \\citet{2009A&A...496..469L}.  From theoretical considerations, both the rapid p-mode as well as the slow g-mode oscillations should be very stable in the pulsating subdwarf B stars (stable in phase, i.e.\\ coherent over time scales of many years, as well as stable in amplitude) due to the identical underlying driving: a $\\kappa$ mechanism acting on the $Z$ bump. Time dependency of phases and amplitudes can however be introduced through nonlinear effects; the resulting mode coupling then starts to invalidate the initial simple assumption. It is therefore crucial to identify target stars that show a stable behaviour.  Under the right conditions, this provides the possibility to obtain meaningful results from long-term photometric monitoring of rapidly pulsating sdB stars: It allows to search for slow changes in pulsation periods due to evolutionary effects, testing the time scales in sdB evolutionary models, and may be used as an additional constraint for sdB asteroseismological models. It also allows to use the star as a clock that provides a regular timing signal, which can be exploited to search for small periodic deviations from the mean pulsation maxima arrival times. Such periodic deviations can be caused by a varying light travel time and may hence indicate a wobble of the pulsating subdwarf B star's location due to one or more low-mass companions, as is the case for \\object{HS\\,2201+2610}. This timing method is sensitive to planetary masses and large period ranges not easily accessible with other methods.  Given the observational efforts required, the detection rate of low-mass companions in wide orbits is surprisingly high so far. Besides timing of pulsations \\citep{2007Natur.449..189S}, timing of the sharp eclipses in HW~Vir type systems has been used to detect tertiary bodies \\citep{2009AJ....137.3181L,2009ApJ...695L.163Q}. Furthermore, \\citet{2009ApJ...702L..96G} have even successfully employed the radial velocity method to detect a close substellar companion to \\object{HD 149382}.  \\begin{table*}[t] %\\begin{center} \\small \\caption{Overview of EXOTIME targets} %% no full stop at the end of caption \\label{tbl:targets} \\begin{tabular}{lllll} \\tableline target             & coordinates (equinox 2000.) & $m_{\\textrm{B}}$ & status & \\\\[1mm] \\tableline {HS\\,2201+2610}    & 22:04:12.0 \\quad +26:25:07 &14.3& {collecting data} & {planet candidate published}, {$\\sin{i}$~unknown}  \\\\[1mm] {HS\\,0702+6043}    & 07:07:09.8 \\quad +60:38:50 &14.7& {collecting data}&{see \\ref{ss:hs0444} and \\ref{s:longterm}} \\\\[1mm] {HS\\,0444+0458}    & 04:47:18.6 \\quad +05:03:35 &15.2 & {collecting data}&{see \\ref{ss:hs0702}} \\\\[1mm] EC\\,09582$-$1137~  & 10:00:41.8 \\quad $-$11:51:35 &15.0& {collecting data}& \\\\[1mm] PG\\,1325+101       & 13:27:48.6 \\quad +09:54:52 &13.8& {collecting data}& \\\\[1mm] \\tableline \\end{tabular} %\\end{center} \\end{table*}  As it has been suggested that the influence of substellar companions or large planetary companions may be decisive for the evolution of subdwarf B stars (see Discussion in section~\\ref{s:discussion}), the EXOTIME program has been set up to find more such objects.  We discuss the target selection in section~\\ref{s:data}.  However, as amplitude variations have repeatedly been reported in sdB pulsators \\citep[e.g.][]{2010ApSS.K}, it is crucial to carefully check the stability of the targets first of all. This contribution mainly discusses the current status of these investigations for two of our objects, \\object{HS\\,0444+0458} and \\object{HS\\,0702+6043} (section~\\ref{s:coherence}).  \\object{HS\\,0444+0458} was first discovered to pulsate by \\citet{2001A&A...378..466O}, and has been further characterised by \\citet{2007MNRAS.378.1049R}. % Rapid oscillations were discovered in \\object{HS\\,0702+6043} by \\citet{2002A&A...386..249D}, simultaneous slow oscillations were reported by \\citet{2006A&A...445L..31S}. The on-going \\mbox{EXOTIME} observations for \\object{HS\\,0702+6043} have also previously been summarised by \\citet{2008CoAst.157..185L,2009CoAst.159...94L}.  %These two and two further objects were initially selected as targets for %the EXOTIME program according to the considerations in the next section. ",
        "conclusions": "\\label{s:summary}  The main result for \\object{HS\\,0444+0458} so far is that the O$-$C diagram for its main pulsation frequency f1 is very well-behaved. The O$-$C diagram for f2 is also approximately flat, but with much larger errors and hence less significance due to the lower S/N. We conclude that the degree of coherence of the pulsations in \\object{HS\\,0444+0458} confirms it as a good target for long-term monitoring. The next few years will show if this reveals a measurable $\\dot{P}$ for the star and/or the signature of an orbiting (planetary) companion.  For \\object{HS\\,0702+6043}, previously thought to show a relatively simple spectrum in the p-mode domain, it turns out (\\citealt{2010ApSS.F}, and the analysis in this work) that a significant amount of work must first be put into developing a very good understanding of the pulsational behaviour before these pulsations can be used as a tool to derive $\\dot{P}$ or even search for companions.    The increasing number of detections of low-mass companions \\citep{2007Natur.449..189S,2009AJ....137.3181L,2009ApJ...695L.163Q,2009ApJ...702L..96G,2010ApSS.G} may indicate the existence of a previously undiscovered population of substellar and planetary bodies around subdwarf B stars.  This has renewed interest in the question, initially raised by \\cite{1998AJ....116.1308S}, whether planets might be relevant for the formation of subdwarf B stars.  %Let's assume here for shortness that the context is well known: %  existence of observed binaries vs. apparently single %  mass spectrum of companions %  associated formation scenarios -> sdB+dM, sdB+WD; NS? BH?  To discuss the possible influence of a substellar companion on the primary star's evolution on the red giant branch, the above systems must be considered individually.  In the case of \\object{HS\\,2201+2610} \\citep{2007Natur.449..189S}, for instance, the detected planet is too far out to have strongly influenced the envelope ejection process. On the other hand, when a star forms planets, it often forms full planetary systems with several members, so that a remaining detectable far-out planet may serve as a tracer of previously present closer-in planets, possibly destroyed during the RGB phase. In this sense, the detection of \\object{V391 Pegasi b} could be interpreted as backing up the suggestion that planets can indeed enhance the mass loss on the RGB in such a way that the subsequent formation of a subdwarf B star is favoured \\citep{2010ApSS.S}. The detection of \\object{HD 149382 b} \\citep{2009ApJ...702L..96G,2010ApSS.G} may constitute more direct evidence, with the actual object having undergone the common envelope phase still present today.  The situation is different for the tertiary bodies in the eclipsing close binary systems \\object{HW Vir} and \\object{HS\\,0705+6700} \\citep{2009AJ....137.3181L,2009ApJ...695L.163Q}. The ''problem'' of ''missing'' companions that would be needed to explain the formation of the primary subdwarf B star through one of the accepted binary evolution channels has never existed here, as suitable secondary bodies were already known.  The detections of low-mass objects around apparently single sdB stars are therefore the more relevant ones when it comes to the discussion whether such companions can be efficient in forming subdwarf B stars. In particular, the amount of angular moment transferred and hence the modification of the ejection efficiency during a common envelope phase are a matter of debate among theorists. %%% (which) Podsiadlowski reference ??? In alternative scenarios, the existence of planets around subdwarf B stars is explained by formation processes leading to second generation planets.  The EXOTIME program has the potential to add to the empirical data available on which such a debate must ultimately be based.  \\begin{comment} 8 pages for talks (minimum 4 pages!) 4 pages per poster SUM 8+4+4= max 16 pages \\end{comment}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{table*}[tb] \\small \\caption{Photometric data archive for HS0444+0458} %% no full stop at the end of caption \\label{tbl:hs0444} \\begin{tabular}{cccrrrlllll} \\tableline  %% rule at top Year & Month & Day & L [h] & N & $\\sigma_{\\textrm{res}}$ & Exp [s] & Site\\tablenotemark{a} & Inst. & Filter & Observer\\tablenotemark{b}\\\\ \\tableline 2008 & Aug & 19 & 0.5 & 107 & 0.0069 & 5 & TNG~3.6m & DOLORES & B & RS, SL\\\\ &     & 21 & 1.0 & 237 & 0.0076 & 5 & TNG~3.6m & DOLORES & B & RS, SL\\\\ &     & 24 & 0.7 & 173 & 0.0060 & 5 & TNG~3.6m & DOLORES & B & RS, SL\\\\ &     & 26 & 1.0 & 224 & 0.0051 & 5 & TNG~3.6m & DOLORES & B & RS, SL\\\\ &     & 27 & 2.1 & 357 & 0.0040 & 10 & TNG~3.6m & DOLORES & B & RS, SL\\\\ &     & 28 & 1.7 & 230 & 0.0042 & 10 & TNG~3.6m & DOLORES & B & RS, SL\\\\ & Sep & 29 & 2.2 & 171 & 0.0089 & 10 & CA~2.2m & CAFOS & B & service (FH)\\\\ & Oct & 24 & 2.8 & 300 & 0.0191 & 12 & LOAO~1.0m & 2k CCD & B & SLK\\\\ &     & 26 & 4.9 & 588 & 0.0232 & 12 & LOAO~1.0m & 2k CCD & B & SLK\\\\ &     & 29 & 1.2 & 163 & 0.0175 & 12 & LOAO~1.0m & 2k CCD & B & SLK\\\\ & Nov & 21 & 3.2 & 263 & 0.0092 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ \\tableline 2009 & Jan & 29 & 2.6 & 275 & 0.0239 & 20 & LOAO~1.0m & 2k CCD & B & SLK\\\\ &     & 31 & 4.7 & 497 & 0.0275 & 20 & LOAO~1.0m & 2k CCD & B & SLK\\\\ & Mar & 20 & 1.0 & 94 & 0.0090 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ &     & 21 & 1.3 & 135 & 0.0121 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ &     & 22 & 1.2 & 124 & 0.0099 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ & Oct & 16 & 1.3 & 195 & 0.0138 & 20 & M/N~1.2m & 1k CCD & B & RL\\\\ &     & 18 & 1.6 & 242 & 0.0148 & 20 & M/N~1.2m & 1k CCD & B & RL\\\\ &     & 22 & 2.1 & 235 & 0.0102 & 10 & CA~2.2m & CAFOS & B & service (AA)\\\\ &     & 23 & 2.9 & 330 & 0.0082 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ &     & 24 & 2.2 & 244 & 0.0075 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ &     & 25 & 1.7 & 104 & 0.0061 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ &     & 26 & 3.4 & 387 & 0.0086 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ \\tableline %% rule at bottom \\end{tabular} % Any table notes must follow the \\end{tabular} command. \\tablenotetext{a}{CA: Calar Alto Observatory, LOAO: Mt. Lemmon Optical Astronomy Observatory, M/N: MONET/North Telescope, TNG: Telescopio Nazionale Galileo} \\tablenotetext{b}{FH: Felipe Hoyo, MA: Manuel Alises, RL: Ronny Lutz, RS: Roberto Silvotti, SLK: Seung-Lee Kim, SL: Silvio Leccia} %\\tablecomments{some more comments here...} \\end{table*}  \\begin{table*}[tb] \\small \\caption{Photometric data archive for HS\\,0702+6043, part a} %% no full stop at the end of caption \\label{tbl:hs0702a} \\begin{tabular}{cccrrrlllll} \\tableline  %% rule at top Year & Month & Day & L [h] & N & $\\sigma_{\\textrm{res}}$ & Exp [s] & Site\\tablenotemark{a} & Inst. & Filter & Observer\\tablenotemark{b}\\\\ \\tableline 1999 & Dec & 06 & 5.0 & 943 & 0.0103 & 10 & CA~1.2m & CCD & - & SD, SS\\\\ &     & 07 & 1.9 & 414 & 0.0141 & ~\\,8 & CA~1.2m & CCD & - & SD, SS\\\\ &     & 08 & 1.4 & 241 & 0.0085 & 10 & CA~1.2m & CCD & - & SD, SS\\\\ \\tableline 2000 & Oct & 08 & 0.7 & 248 & 0.0027 & 10 & NOT~2.56m & ALFOSC & 3W & R\\O\\\\ \\tableline 2004 & Feb & 04 & 7.3 & 685 & 0.0233 & 30 & Tue~0.8m & ST7 & G & TN \\\\ &     & 09 & 5.8 & 244 & 0.0016 & 60 & StB~2.2m & 2k CCD & F555W & EMG\\\\ &     & 10 & 6.2 & 264 & 0.0014 & 60 & StB~2.2m & 2k CCD & F555W & EMG\\\\ \\tableline 2005 & Jan & 17-27 & 56.0 & 5981 & 0.0082 & 20, 30 & CA~2.2m & CAFOS & B & SS, TS\\\\ \\tableline %2007-08 & Nov-Mar &  & 424.0 & 31638 & 0.0027 &  & MtB~1.55m & MONT~4k & S 8612 & EMG\\\\ 2007 & Dec & 18 & 5.9 & 498 & 0.0442 & 40 & Tue~0.8m & SBIG ST-L & B & TN\\\\ &     & 19 & 6.7 & 557 & 0.0589 & 40 & Tue~0.8m & SBIG ST-L & B & TN\\\\ &     & 20 & 10.0 & 842 & 0.0994 & 40 & Tue~0.8m & SBIG ST-L & B & TN\\\\ &     & 21 & 9.3 & 760 & 0.1350 & 40 & Tue~0.8m & SBIG ST-L & B & TN\\\\ &     &    & 12.1 & 621 & 0.0200 &  & Goe~0.5m & SBIG ST-L & B & SS, RL, RK\\\\ \\tableline 2007 & Nov &    &    &    &    &    &    &    &    &\\\\[-1ex] \\vdots   &  \\vdots  &    & 424.0 & 31638 & 0.0027 & 50 & MtB~1.55m & MONT~4k & S 8612 & EMG \\& coll.\\\\ 2008 & Mar &    &    &    &    &    &    &    &    &\\\\ \\tableline 2008 & Feb & 08 & 7.2 & 601 & 0.0124 & 40 & Tue~0.8m & SBIG ST-L & B & TN\\\\ &     & 09 & 7.1 & 590 & 0.0121 & 40 & Tue~0.8m & SBIG ST-L & B & TN\\\\ &     &   & 8.9 & 642 & 0.0381 & 35 & Goe~0.5m & SBIG ST-L & B & RL, BB\\\\ &     & 10 & 5.8 & 488 & 0.0134 & 40 & Tue~0.8m & SBIG ST-L & B & TN\\\\ &     &    & 12.2 & 910 & 0.0324 & 35 & Goe~0.5m & SBIG ST-L & B & RL, MH\\\\ &     & 11 & 4.5 & 338 & 0.0439 & 35 & Goe~0.5m & SBIG ST-L & B & RL, MH\\\\ &     & 13 & 1.1 & 86 & 0.0749 & 35 & Goe~0.5m & SBIG ST-L & B & RL\\\\ &     & 15 & 5.4 & 379 & 0.0552 & 35 & Goe~0.5m & SBIG ST-L & B & RL, SW\\\\ &     & 29 & 3.7 & 360 &  & 20 & Loi~1.5m & & B  & service (ADB)\\\\ & Mar & 01 & 5.7 & 451 &  & 25 & Loi~1.5m & & B  & service (ADB)\\\\ &     & 03 & 5.1 & 429 &  & 25 & Loi~1.5m & & B  & service (ADB)\\\\ &     & 05 & 0.6 & 47 & 0.0238 & 35 & Goe~0.5m & SBIG ST-L & B & RL, SS\\\\ &     & 13 & 3.7 & 126 & 0.0069 & 30 & Asi~1.8 & Afosc & R & SB, RC\\\\ &     &    & 2.3 & 248 & 0.0105 & 30 & Kon~1.0m & PI VA 1300B & B & MP, PP\\\\ &     & 15 & 4.6 & 597 & 0.0095 & 25 & Kon~1.0m & PI VA 1300B & B & MP, PP\\\\ &     & 18 & 3.7 & 348 & 0.0111 & 35 & Kon~1.0m & PI VA 1300B & B & MP, PP\\\\ &     & 19 & 1.1 & 137 & 0.0088 & 25 & Kon~1.0m & PI VA 1300B & B & MP, PP\\\\ & May & 11 & 3.0 & 204 & 0.0315 & 40 & Goe~0.5m & SBIG ST-L & B & SW\\\\ &     & 12 & 1.2 & 160 & 0.0380 & 40 & Goe~0.5m & SBIG ST-L & B & SW\\\\ &     & 13 & 3.2 & 266 & 0.0254 & 40 & Tue~0.8m & SBIG ST-L & B & AH\\\\ & Oct & 15 & 1.3 & 113 & 0.0085 & 10 & CA~2.2m & CAFOS & B & service (LM)\\\\ &     & 16 & 0.8 & 65  & 0.0101 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ &     & 17 & 0.5 & 48  & 0.0152 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ &     & 18 & 1.6 & 143 & 0.0068 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ &     & 25 & 2.6 & 285 & 0.0089 & 30 & Kon~1.0m & PI VA 1300B & B & ZB\\\\ &     &    & 4.4 & 518 & 0.0107 & 15 & LOAO~1.0m & 2k CCD & B & SLK\\\\ &     & 28 & 2.9 & 326 & 0.0104 & 15 & LOAO~1.0m & 2k CCD & B & SLK\\\\ &     & 30 & 3.9 & 455 & 0.0104 & 15 & LOAO~1.0m & 2k CCD & B & SLK\\\\ & Nov & 21 & 3.0 & 254 & 0.0048 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ &     & 23 & 5.7 & 496 & 0.0064 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ & Dec & 16 & 3.0 & 243 & 0.0068 & 25 & M/N~1.2m & 1k CCD & B & RL, BL\\\\ &     & 20 & 2.7 & 202 & 0.0116 & 25 & M/N~1.2m & 1k CCD & B & RL\\\\ \\tableline 2009 & Jan & 06 & 1.6 & 143 & 0.0364 & 30 & Goe~0.5m & SBIG ST-L & B & RL, US\\\\ &     & 22 & 3.7 & 334 & 0.0080 & 25 & M/N~1.2m & 1k CCD & B & RL\\\\ &     & 28 & 5.4 & 473 & 0.0147 & 20 & LOAO~1.0m & 2k CCD & B & SLK\\\\ &     & 30 & 6.6 & 662 & 0.0138 & 20 & LOAO~1.0m & 2k CCD & B & SLK\\\\ \\tableline %% rule at bottom \\end{tabular} % Any table notes must follow the \\end{tabular} command. \\tablenotetext{a}{see footnote in Table \\ref{tbl:hs0702b} for the site abbreviations} \\tablenotetext{b}{see footnote in Table \\ref{tbl:hs0702b} for the observer abbreviations} %\\tablecomments{more comments here...} \\end{table*}  \\begin{table*}[tb] \\small \\caption{Photometric data archive for HS\\,0702+6043, part b} %% no full stop at the end of caption \\label{tbl:hs0702b} \\begin{tabular}{cccrrrlllll} \\tableline  %% rule at top Year & Month & Day & L [h] & N & $\\sigma_{\\textrm{res}}$ & Exp [s] & Site\\tablenotemark{a} & Inst. & Filter & Observer\\tablenotemark{b}\\\\ \\tableline 2009 & Feb & 18 & 2.4 & 202 & 0.0256 & 30 & Goe~0.5m & SBIG ST-L & B & RL, TOH\\\\ %     &     & 23-28 &  &  &  &  & TT1~1.5m &  &  & DM, MD, FC\\\\ &     & 23 & 3.6 & 304 & 0.0126 & 25 & M/N~1.2m & 1k CCD & B & RL, US, MH\\\\ &     & 25 & 6.4 & 564 & 0.0076 & 25 & M/N~1.2m & 1k CCD & B & RL\\\\ &     & 26 & 6.5 & 536 & 0.0111 & 25 & M/N~1.2m & 1k CCD & B & RL, LN\\\\ &     & 27 & 6.6 & 556 & 0.0083 & 25 & M/N~1.2m & 1k CCD & B & RL\\\\ &     &    & 5.7 & 532 & 0.0067 & 15 & Asi~1.8m & Afosc & no & SB, RC\\\\ & Mar & 01 & 6.2 & 497 & 0.0090 & 25 & M/N~1.2m & 1k CCD & B & RL\\\\ &     & 18 & 3.3 & 279 & 0.0060 & 25 & M/N~1.2m & 1k CCD & B & RL\\\\ &     & 19 & 3.6 & 291 & 0.0098 & 25 & M/N~1.2m & 1k CCD & B & RL\\\\ &     & 20 & 2.4 & 243 & 0.0058 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ &     & 21 & 2.4 & 232 & 0.0068 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ &     &    & 6.0 & 351 & 0.0150 & 30 & Asi~1.8m & Afosc & B & SB, RC\\\\ &     & 22 & 2.4 & 245 & 0.0049 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ &     &    & 4.9 & 279 & 0.0140 & 30 & Asi~1.8m & Afosc & B & SB, RC\\\\ &     & 24 & 3.4 & 243 & 0.0092 & 25 & M/N~1.2m & 1k CCD & B & RL\\\\ &     & 26 & 2.8 & 251 & 0.0081 & 25 & M/N~1.2m & 1k CCD & B & RL, US\\\\ &     &    & 3.9 & 396 & 0.0118 & 20 & LOAO~1.0m & 2k CCD & B & SLK\\\\ &     & 27 & 1.3 & 160 & 0.0185 & 20 & LOAO~1.0m & 2k CCD & B & SLK\\\\ &     & 28 & 3.9 & 404 & 0.0172 & 20 & LOAO~1.0m & 2k CCD & B & SLK\\\\ &     & 29 & 2.9 & 309 & 0.0112 & 20 & LOAO~1.0m & 2k CCD & B & SLK\\\\ &     & 30 & 3.8 & 386 & 0.0144 & 20 & LOAO~1.0m & 2k CCD & B & SLK\\\\ &     & 31 & 4.0 & 399 & 0.0154 & 20 & LOAO~1.0m & 2k CCD & B & SLK\\\\ & Sep & 27 & 3.9 & 406 & 0.0132 & 25 & M/N~1.2m & 1k CCD & B & RL\\\\ &     & 28 & 4.0 & 420 & 0.0117 & 25 & M/N~1.2m & 1k CCD & B & RL\\\\ & Oct & 16 & 4.0 & 412 & 0.0099 & 25 & M/N~1.2m & 1k CCD & B & RL\\\\ &     & 18 & 0.9 & 100 & 0.0109 & 25 & M/N~1.2m & 1k CCD & B & RL\\\\ &     & 19 & 6.1 & 645 & 0.0119 & 25 & M/N~1.2m & 1k CCD & B & RL\\\\ &     & 22 & 3.9 & 359 & 0.0076 & 10 & CA~2.2m & CAFOS & B & service (AA)\\\\ &     & 23 & 3.6 & 334 & 0.0067 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ &     & 24 & 3.1 & 289 & 0.0061 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ &     & 25 & 2.8 & 297 & 0.0082 & 25 & M/N~1.2m & 1k CCD & B & RL\\\\ &     &    & 3.4 & 315 & 0.0078 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ &     & 26 & 3.0 & 279 & 0.0064 & 10 & CA~2.2m & CAFOS & B & service (MA)\\\\ \\tableline %% rule at bottom \\end{tabular} % Any table notes must follow the \\end{tabular} command. \\tablenotetext{a}{Asi: Asiago 182cm Copernico Telescope, CA: Calar Alto Observatory, Goe: G\\\"ottingen IAG 50cm Telescope, Kon: Konkoly RCC Telescope, LOAO: Mt. Lemmon Optical Astronomy Observatory, Loi: Loiano 152cm Telescope, M/N: MONET/North Telescope, MtB: Mt. Bigelow Kuiper Telescope, NOT: Nordic Optical Telescope, StB: Steward Observatory Bok Telescope, Tue: T\\\"ubingen 80cm Telescope} \\tablenotetext{b}{AA: Alberto Aguirre, ADB: Antonio De Blasi, AH: Agnes Hoffmann, BB: Benjamin Beeck, BL: Bj\\\"orn Loeptien, EMG: Elizabeth M.~Green, LM: Luzma Montoya, LN: Lisa Nortmann, MA: Manuel Alises, MH: Markus Hundertmark, MP: Margit Papar\\'o, PP: P\\'eter P\\'apics, RC: Riccardo Claudi, RK: Renate Kruspe, RL: Ronny Lutz, R\\O: Roy \\O stensen, SB: Serena Benatti, SD: Stefan Dreizler, SLK: Seung-Lee Kim, SS: Sonja Schuh, SW: Sascha Werhahn, TN: Thorsten Nagel, TOH: Tim-Oliver Husser, TS: Thorsten Stahn, US: Ulf Seemann, ZB: Zs\\'ofia Bogn\\'ar} %\\tablecomments{comments here...} \\end{table*}   %% Acknowledgements % %"
    },
    "1005/1005.1663_arXiv.txt": {
        "abstract": "Knowledge of an exoplanet's oblateness and obliquity would give clues about its formation and internal structure. In principle, a light curve of a transiting planet bears information about the planet's shape, but previous work has shown that the oblateness-induced signal will be extremely difficult to detect. Here we investigate the potentially larger signals due to planetary spin precession. The most readily detectable effects are transit depth variations (\\tdv) in a sequence of light curves. For a planet as oblate as Jupiter or Saturn, the transit depth will undergo fractional variations of order $1\\%$. The most promising systems are those with orbital periods of approximately 15--30 days, which is short enough for the precession period to be less than about 40 years, and long enough to avoid spin-down due to tidal friction. The detectability of the T$\\delta$V signal would be enhanced by moons (which would decrease the precession period) or planetary rings (which would increase the amplitude). The {\\em Kepler} mission should find several planets for which precession-induced T$\\delta$V signals will be detectable.  Due to modeling degeneracies, {\\em Kepler} photometry would yield only a lower bound on oblateness.  The degeneracy could be lifted by observing the oblateness-induced asymmetry in at least one transit light curve, or by making assumptions about the planetary interior. ",
        "introduction": "Measuring the oblateness of exoplanets would further our understanding of planetary formation, rotation, and internal structure. One possible measurement technique relies on the differences between the transit light curve of a spherical planet and an oblate planet with the same sky-projected area (Seager \\& Hui 2002, Barnes \\& Fortney 2003). However, the differences are minuscule, of order 200 parts per million (ppm) for a planet as oblate as Saturn, and 2~ppm for a ``hot Jupiter'' whose spin rate has been slowed by tidal friction into synchronization with its orbit. In a previous paper we showed that with this technique, even the best available light curves are only barely sufficient to rule out a Saturn-like oblateness (Carter \\& Winn 2009). Those results pertained to the planet HD~189733b, for which one would expect spin-orbit synchronization, and consequently the theoretical oblateness was an order of magnitude below the empirical upper limit.  In that work we also pointed out the potentially observable effects of a phenomenon that had been previously overlooked: the precession of the planet's rotation axis. Precession of an oblate planet causes the sky-projected area of the planet to change over time, thereby causing gradual changes in the depth and duration of transits. In this paper we investigate the observable manifestations of spin precession in a broader context. Section 2 describes the characteristics of the signal in terms of the properties of the star, planet, and orbit. Section 3 discusses the range of orbital periods for which the signal is most readily detectable, considering the timescales for precession and spin-orbit synchronization. Section 4 presents simulated results for a specific case, a Saturn-like planet observed by the {\\it Kepler} satellite. Section 5 summarizes and discusses the results. \\\\ \\\\ ",
        "conclusions": "\\label{sec:disc}  In this paper we have investigated the observability of changes to transit light curves resulting from the spin precession of an oblate, oblique exoplanet.  The most readily detectable signal is the \\tdv signal, the variation in transit depth due to the changing area of the planetary silhouette.  The planets that seem most likely to exhibit detectable effects are those with periods between 15--30 days (around Sun-like stars), which is short enough for precession periods to be 40~yr or less, and long enough to hope that tidal spin-orbit synchronization has not taken place.  It is also important to consider other physical processes that could give rise to T$\\delta$V signals, and which might confound the interpretation of the data. Starspots and other types of stellar variability can produce transit depth variations. These can be recognized and taken into account by monitoring the star outside of transits, as is done automatically by the {\\it Kepler} satellite. We are not aware of any atmospheric phenomena associated with the exoplanet that would result in variations in the projected area at the 1\\% level, mimicking those due to uniform spin precession.  However, we can think of two plausible phenomena that would affect the T$\\delta$V signal: moons and rings. If the planet has any moons then the precession period may be shorter than we have calculated. Satellites provide more leverage for the star to torque the exoplanetary system.  Mathematically, satellites augment the effective values of $J_2$ and $\\mathds{C}$. Following Ward \\& Hamilton (2004) we may write the enhanced values as $J_2+j$ and $\\mathds{C}+c$, where \\begin{eqnarray} c & = & \\sum_i \\frac{m_i}{M_p} \\left(\\frac{a_i}{R_{\\rm eq}}\\right)^2 \\frac{P_{\\rm rot}}{p^i_{\\rm orb}}\\;\\;\\;~\\mbox{and}\\\\ j & = &\\frac{1}{2} \\sum_i \\frac{m_i}{M_p} \\left(\\frac{a_i}{R_{\\rm eq}}\\right)^2 \\frac{\\sin \\left(\\theta-I_i\\right)}{\\sin \\theta} \\end{eqnarray} where $m_i$, $a_i$, $I_i$ and $p^i_{\\rm orb}$ are the satellites' masses, orbital radii, orbital inclinations (relative to the planetary equator), and orbital periods.  For example, in the Saturnian system $j/J_2 \\approx 3.2$ while $c/\\mathds{C} \\approx 0.01$ such that $(\\mathds{C}+c)/(J_2+j) \\approx 3.2$ (as compared to $\\mathds{C}/J_2 = 13.5$).  Titan alone is responsible for about $90\\%$ of $j$ and $c$, shortening Saturn's precession period by a factor of four relative to a satellite-free Saturn.\\footnote{For Titan: $a_i/R_{\\rm eq} = 20.2$, $m_i/M_p = 2.3\\times10^{-4}$, $p^i_{\\rm orb}/P_{\\rm rot}=38.3$; Murray \\& Dermott (2000).} Fig.~(\\ref{fig:times}) shows the effect of Titan analogs at various distances around the Saturn analog considered in this paper. The distances are expressed as fractions of the Hill radius $r_H$. This effect might be used to implicate the presence of exomoons, if a T$\\delta$V signal were observed and found to correspond to an effective $J_2$ too large to be plausibly attributed to the planet alone.  Ring systems would increase the amplitude of the signal, while leaving the period unchanged. Optically thick rings that lie within the equatorial plane of the planet would increase the fractional variation of the sky-projected area, as the planet precesses. In contrast to exomoons, ring systems have little mass, and would not significantly reduce the precessional period. This is the situation in the Saturnian system (Ward \\& Hamilton 2004). It is not certain that rings could exist around planets having orbital periods between 15 and 30 days. Rings around planets with orbital periods less than $P_{\\rm orb} \\approx15$ days would likely be short-lived as a result of Poynting-Roberston drag amongst other destructive effects (Barnes \\& Fortney 2004). Also, we would be unlikely to find rings composed of water ice (e.g.\\ Saturn's rings) around planets whose orbits are interior to the snow line ($< 1$ AU). Nevertheless, rings of other compositions may exist.  In short, rings and moons would each affect the observed T$\\delta$V signal, and in complementary ways. This may introduce some ambiguity in the estimation of the shape parameters of the rotating planet, but may also allow the rings and moons to be detectable, issues that we leave for future work.  For simplicity we have considered only circular orbits.  Planets on eccentric orbits will undergo apsidal precession and nodal precession, which will result in time variable stellar torques on the planet and consequent modifications to the spin-axis precession. For Saturn or Jupiter analogs at $P_{\\rm orb} \\gtrsim 15$ days, the apsidal and nodal precession periods are $\\gtrsim$$10^7$~yr and therefore likely to be irrelevant (Ragozzine \\& Wolf 2009).  In addition, we restricted our attention to the simplest case of uniform spin precession, but in reality the perturbations from other bodies may cause the spin axis to perform a more complex ballet. For example, Mars's spin axis tumbles chaotically (Touma \\& Wisdom 1992). Saturn's moon Titan causes a 700~yr modulation of Saturn's spin precession frequency, due to its inclined orbit. Furthermore, Saturn's spin axis may be trapped in a resonance with Neptune's orbit, causing it to librate with an angular amplitude of $\\gtrsim$$31^\\circ$ as it circulates about the second Cassini state (Ward \\& Hamilton 2004). These effects would be manifested as additional time dependences (``noise'') in the \\tdv signal. The effects are impossible to forecast for exoplanets, depending as they do on the existence of other bodies and any resonances that may occur."
    }
}